{
    "1107/1107.1239_arXiv.txt": {
        "abstract": "Mounting discoveries of extrasolar planets orbiting post-main sequence stars motivate studies aimed at understanding the fate of these planets.  In the traditional ``adiabatic\" approximation, a secondary's eccentricity remains constant during stellar mass loss. Here, we remove this approximation, investigate the full two-body point-mass problem with isotropic mass loss, and illustrate the resulting dynamical evolution.  The magnitude and duration of a star's mass loss combined with a secondary's initial orbital characteristics might provoke ejection, modest eccentricity pumping, or even circularisation of the orbit.  We conclude that Oort clouds and wide-separation planets may be dynamically ejected from $1 M_{\\odot}-7 M_{\\odot}$ parent stars during AGB evolution. The vast majority of planetary material which survives a supernova from a $7 M_{\\odot}-20 M_{\\odot}$ progenitor will be dynamically ejected from the system, placing limits on the existence of first-generation pulsar planets. Planets around $>20 M_{\\odot}$ black hole progenitors may easily survive or readily be ejected depending on the core collapse and superwind models applied. Material ejected during stellar evolution might contribute significantly to the free-floating planetary population. ",
        "introduction": "Understanding the formation and subsequent dynamical evolution of exoplanets has been a motivational hallmark for many observational and theoretical investigations.  However, extrasolar planets continue to be discovered in surprising and exotic environments, and questions about the {\\it endstate} of exoplanets are becoming increasingly relevant. Few studies so far have modeled these systems, which often feature evolved and variable parent stars. The rich dynamics therein fundamentally differ from studies of planets around main sequence stars.  Examples of exoplanets which do not orbit main sequence stars are growing.  The first confirmed extrasolar planets were discovered around a neutron star: specifically, the millisecond pulsar PSR1257+12 \\citep{wolfra1992,wolszczan1994}. The minimum masses of these three planets continue to be among the lowest known to date, and two of these planets resonantly interact.  \\cite{sigurdsson2003} later discovered another pulsar planet, around the binary radio millisecond pulsar PSR B1620-26. Exoplanets are also thought to orbit white dwarfs and stars with white dwarf companions.  In the first category, GD 66 \\citep{muletal2008,muletal2009}, GD 356 \\citep{wicetal2010} and Gliese 3483 (Matt Burleigh, private communication) are planet-hosting stars.  In the second category, examples are thought to include Gl 86 $=$ HD 13445 \\citep{queetal2000,mugneu2005,lagetal2006}, HD 27442 \\citep{butetal2001,chaetal2006}, and HD 147513 \\citep{mayetal2004,desbar2007}.  Additionally, planets have been discovered orbiting stars that have turned off of the main sequence but are not yet stellar remnants.  \\cite{siletal2007} discovered a giant planet orbiting the extreme horizontal branch star V 391 Pegasi, \\cite{geietal2009} found a planet around the hot subdwarf star HD 149382, \\cite{leeetal2009} reported circumbinary planets to the sdB+M eclipsing system HW Virginis, and \\cite{setetal2010} suggested that the planet orbiting the red horizontal branch star HIP 13044b might be of extragalactic origin.  Cataclysmic variables are another class of systems which might harbor planets, and recently, planets around the cataclysmic variables QS Vir \\citep{qianetal2010a}, DP Leo \\citep{qianetal2010b} and HU Aqr \\citep{qiaetal2011} have been reported.  Prospects for discovering additional planets orbiting white dwarfs \\citep{draetal2010,faeetal2011} and extreme horizontal branch stars \\citep{beasok2011} are promising, and observational campaigns to do so have already been initiated \\citep{hogetal2009,benetal2010,schetal2010}.  The {\\it Kepler} mission can detect even smaller bodies around white dwarfs \\citep{disetal2010}.  Theoretical investigations regarding the evolution of planets around post-main sequence stars have focused primarily on planet engulfment and interaction with the expanding stellar envelope, both for exoplanets and specifically for the Earth. \\cite{villiv2007}, \\cite{massarotti2008} and \\cite{villiv2009} use particular stellar evolutionary tracks to determine ranges of semimajor axes at which planets are likely to be engulfed. In this regime, tidal modelling has a significant effect on the subsequent orbital evolution.  However, as summarized by \\cite{hansen2010}, the nature of tidal dissipation is poorly understood and continues to yield different results depending on the model and assumptions used.  For this reason, the fate of the Earth is uncertain. \\cite{sacetal1993}, \\cite{rybden2001}, \\cite{schcon2008} and \\cite{iorio2010} all explore the fate of the Earth in light of the Sun's post main-sequence mass loss, with differing results.  Alternatively, \\cite{debsig2002} focus on the stability of multi-planet systems and link stellar mass loss to instability timescales. By doing so, they demonstrate how multiple planets beyond the reach of the star's expanding envelope might become unstable.  In this study, we consider just a single planet, or smaller body. We perform a detailed analysis of the variable mass two-body problem and apply the results to a wide range of star-planet fates that encompass all stellar masses $\\lesssim 150 M_{\\odot}$.  We focus on how stellar mass loss affects the eccentricity of a planet or planetary material, a link often ignored in previous studies.  As a result, we show that planetary material can be ejected from a system based on mass loss alone.  We then quantify for what combination of parameters we can expect this behavior.  We start, in Section 2, by reviewing the history of the variable mass two-body problem and the corresponding equations of motion.  We then analyze the orbital evolution in different mass loss regimes, determine where and when the traditionally-used adiabatic approximation holds, and estimate when the planets would become unstable.  In Section 3, we apply the theory to stars of all masses up to $150M_{\\odot}$ in order to pinpoint realistic systems which would yield instability.  We treat five different mass regimes in separate subsections.  We then discuss the caveats, implications and potential extensions in Section 4, and conclude in Section 5. ",
        "conclusions": "The variable-mass two-body problem allows for the bodies to become unbound or highly eccentric.  The implications of this physical principle affect all dying stellar systems which contain any orbiting material.  Many Oort clouds and wide-orbit planets will have their orbits disrupted. The extent of the disruption depends crucially on their initial semimajor axes, eccentricities, and true anomalies, and the subtleties of stellar evolution.  Stars with progenitor masses of $4 M_{\\odot}-8 M_{\\odot}$ will readily eject objects that are beyond a few hundred AU distant, and excite the eccentricities of the remaining bound material at that distance. Supernovae which produce neutron stars eject nearly but not all orbiting material.  Conversely, other exotic systems, such as those with black holes, could have easily retained planets during their formation. Stellar mass loss might be the dominant source of the free-floating planet population, and orbital properties of currently observed disrupted planets in aged systems may be tracers of the evolution of their parent stars."
    },
    "1107/1107.1713_arXiv.txt": {
        "abstract": "We present a description of the CL-based package {\\sc xdspres}, which aims at being a complete reducing facility for cross-dispersed spectra taken with the Ohio State Infrared Imager/Spectrometer, as installed at the SOAR telescope. This instrument provides spectra in the range between 1.2$\\upmu$m and 2.35$\\upmu$m in a single exposure, with resolving power of R $\\sim$ 1200. {\\sc xdspres} consists of two tasks, namely \\textit{xdflat} and \\textit{doosiris}. The former is a completely automated code for preparing normalized flat field images from raw flat field exposures. \\textit{Doosiris} was designed to be a complete reduction pipeline, requiring a minimum of user interaction. General steps towards a fully reduced spectrum are explained, as well as the approach adopted by our code. The software is available to the community through the web site \\textit{http://www.if.ufrgs.br/$\\sim$ruschel/software}. ",
        "introduction": "Cross-dispersed spectroscopy makes possible to acquire information of wide spectral regions in a single exposure, by projecting several dispersion axes on the detector simultaneously. As a consequence, the reduction process required to analyze this kind of data is complicated, since different diffraction orders need to be selected, extracted, calibrated independently and combined in the final step. This difficulty led many authors to develop methods and software packages for the reduction of cross-dispersed and echelle spectra \\citep[e.g.][]{moreno1982, rossi1985, piskunov2002, bochanski2009}.  In the past decade the near infrared (NIR) has also been explored by cross-dispersed spectrographs, such as Spex \\citep{rayner2003} at the NASA Infrared Telescope Facility (IRTF), with a resolving power of $\\sim$ 2000 and reaching from 0.8 to 5.5$\\upmu$m. Other examples are TripleSpec \\citep{edelstein2007} and the Folded-port Infrared Echellette (FIRE) \\citep{simcoe2008}, achieving R $\\sim$ 2600 and R $\\sim$ 6000 respectively, and covering roughly the same wavelength domain (0.8 - 2.4$\\upmu$m).  Another instrument of similar capabilities is the Ohio State Infrared Imager/Spectrometer (OSIRIS), currently installed at the Southern Astrophysics Research Observatory (SOAR), attached to the 4.1m telescope. OSIRIS provides spectral coverage from 1.0$\\upmu$m to 2.4$\\upmu$m in cross-dispersed mode, with a resolving power of $\\sim$ 1200. High resolution (R $\\sim$ 3000) long-slit modes are also available, but multi-band spectroscopy of this kind suffers from  differences in aperture and seeing.  However, reduction of NIR spectra has a complexity of its own, mostly related to telluric spectral features, both in absorption and emission, and black body radiation due to the telescope itself. A rich literature has been developed on the subject \\citep[e.g.][]{maiolino1996,vacca2003,cushing2004}.  There are currently no specific software packages available for the reduction of cross-dispersed spectra taken with OSIRIS. Aiming at providing a fast and highly automated task, we developed the {\\sc xdspres} (acronym for cross-dispersed spectra reduction script) package. The CL language was chosen due to the availability of almost all of the basic tasks needed to perform the reduction in the Image Reduction and Analysis Facility ({\\sc IRAF}) software \\citep{IRAF1,IRAF2}.  In \\S \\ref{sec:osiris} we describe main aspects of the instrument, focusing on its effects on the reduction process. In \\S \\ref{sec:reduction} we describe the general steps towards a fully reduced spectrum, as well as the approach adopted by the {\\sc xdspres} package to each of these steps, and finally in \\S \\ref{sec:summary} we give a brief summary. ",
        "conclusions": "\\label{sec:summary}  We have presented the {\\sc xdspres} CL-based package, consisting of the \\textit{xdflat} and \\textit{doosiris} tasks, aimed at being a complete reduction facility for cross-dispersed spectra taken with the OSIRIS spectrometer, currently installed at the SOAR telescope. This particular instrument provides a relatively large spectral coverage, being able to project the full range between 1.2$\\upmu$m and 2.35$\\upmu$m over the detector in a single exposure. The blazing of different orders in the same image adds complexity to the already lengthy reduction of infrared spectroscopy data. {\\sc xdspres} automatically performs the more mechanical and time consuming steps of the reduction, at the same time that it allows considerable user interaction in the more subjective stages. In addition, the possibility of a fast reduction provides means to make site adjustments to the observation strategy. As a sample of actually published data that was fully reduced with the {\\sc xdspres} tasks, see \\citet{riffel2011}. The complete software package and its documentation is available to the community at the web site \\textit{http://www.if.ufrgs.br/$\\sim$ruschel/software}.  \\subsection*"
    },
    "1107/1107.5240_arXiv.txt": {
        "abstract": "{We present two new aspects of Extensive Air Shower (EAS) development universality allowing to make accurate estimation of muon and electromagnetic (EM) shower contents in two independent ways. In the first case, to get muon (or EM) signal in water Cherenkov detectors it is enough to know the vertical depth of shower maximum and the total signal. In the second case, the EM signal can be calculated from the primary particle energy and the zenith angle. In both cases the parameterizations of muon and EM signals are almost independent on primary particle nature, energy and zenith angle.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.0186_arXiv.txt": {
        "abstract": "{``High-resolution'', or ``long-baseline'', science with the SKA and its precursors covers a broad range of topics in astrophysics. In several research areas, the coupling between improved brightness sensitivity of the SKA and a sub-arcsecond resolution would uncover truly unique avenues and opportunities for studying extreme states of matter, vicinity of compact relativistic objects, and complex processes in astrophysical plasmas. At the same time, long baselines would secure excellent positional and astrometric measurements with the SKA and critically enhance SKA image fidelity at all scales. The latter aspect may also have a substantial impact on the survey speed of the SKA, thus affecting several key science projects of the instrument.} ",
        "introduction": "\\label{sec:1}  The benchmark design for the SKA Phase 1 \\cite{schilizzi2007}, envisaging operations in the 0.3-10\\,GHz range and on baselines of up to several hundred kilometres, would have enabled addressing a range of scientific areas relying on sub-arcsecond resolution, including astrometry, pulsar proper motions, supernovae, astrophysical masers, nuclear regions of AGN, physics of relativistic and mildly relativistic outflows, kinetic feedback from AGN, evolution of supermassive black holes and their host galaxies \\cite{carilli2004}. The revised specifications for the SKA$_1$ \\cite{dewdney2010,garrett2010}, shifting the operational frequency range to 0.07-3\\,GHz and limiting the baseline length to 100\\,km, leads to a reduction of the instrumental resolution to 0{\\farcs}3--1{\\farcs}4 for the dishes (in the 0.45--3\\,GHz range) and 1{\\arcs}--8{\\arcs} for the aperture array (in the 0.07--0.45\\,GHz range).  The ``stand alone'' resolution of SKA$_1$ will therefore be sufficient for addressing only a subset of topics listed above. Achieving a higher resolution would rely on inclusion of external antennas and operating in the VLBI mode. This would be mostly feasible for the dish part of SKA$_1$, as most of the present day VLBI arrays are operating at frequencies above 600\\,MHz, and there are no definite plans to extend VLBI operations to below 300\\,MHz.  SKA$_1$ operating in Australia can be an integral part of the LBA/NZ Network and EAVN. It would also have a somewhat limited common visibility with the VLBA.  SKA$_1$ sited in South Africa will be a natural partner to EVN+ antennas. In both cases, collaboration with geodetic VLBI is possible, if 2.3\\,GHz will be maintained as a network frequency by the IVS.  In addition, the sub-arcsecond resolution of SKA$_1$ may actually be an essential requirement also for achieving the specifications envisaged for the traditional ``low-resolution'' science, including the surveying capabilities of the array (often viewed as a backbone of the instrument).  These two aspects of the relevance of long baselines to achieving the scientific goals of SKA$_1$ are discussed below, with Section~\\ref{apl-sec2} describing potential areas of broad scientific impact of high-resolution studies with SKA$_1$ and SKA Precursors and Section~\\ref{apl-sec3} discussing the effect of long baselines on the quality of imaging and surveying capability of SKA$_1$. ",
        "conclusions": "SKA$_1$ would be able to address a number of important astrophysical areas of study relying on high-resolution radio observations: studying supernovae; providing a good account of starburst activity in galaxies; using megamasers and nuclear absorption to probe the nuclear gas in galaxies; understanding in detail the physics of (ultra- and mildly-relativistic) outflows and their connection to the nuclear regions in galaxies; searching for radio emission from weaker AGN and secondary black holes in post-merger galaxies; and addressing the questions of relic activity and activity cycles in AGN.  Reliable high resolution imaging is also needed for achieving the scientific goals in the key science areas of SKA$_1$ requiring low-resolution imaging and surveying of large areas in the sky, most notably the science areas considered as the very ``selling point'' of the instrument."
    },
    "1107/1107.0465_arXiv.txt": {
        "abstract": "The INTEGRAL satellite, observing the sky at high energy, has quadrupled the number of supergiant X-ray Binaries known in the Galaxy and has revealed new populations of previously hidden High Mass X-ray Binaries. These observations raise new questions about the formation and evolution of these sources. The number of detected sources is now high enough to allow us to carry out a statistical analysis of the distribution of HMXBs in the Milky Way. We derive the distance of each HMXB using a Spectral Energy Distribution fitting procedure, and we examine the correlation with the distribution of star forming complexes (SFCs) in the Galaxy. We show that HMXBs are clustered with SFCs, with a typical size of 0.3 kpc and a characteristic distance between clusters of 1.7 kpc. ",
        "introduction": "High Mass X-ray Binaries (HMXBs) are binary systems composed of a compact object, a neutron star or a black hole candidate, accreting matter from a massive companion star: either a main sequence Be star or an evolved supergiant O or B star. Most of these sources are observed in the Galactic Plane \\citep{Bird_2007} as it is expected for such young star systems which do not have time to move far from their birthplaces.\\\\ Thanks to the dedicated observations from \\textit{RXTE} and \\textit{INTEGRAL}, around 200 HMXBs are currently known in the Milky Way allowing us to focus on their distribution. Using RXTE data, \\citet{Grimm_2002} highlighted clear signatures of the spiral structure in the spatial distribution of HMXBs. In the same way, \\citet{Dean_2005}, \\citet{Lutovinov_2005}, \\citet{Bodaghee_2007} and \\citet{Bodaghee_2011} showed that HMXBs observed with \\textit{INTEGRAL} also seem to be associated with the spiral structure of the Galaxy. However, the HMXB positions, mostly derived from their X-ray luminosity, are not well constrained and highly uncertain due to direct accretion as in HMXB. In order to overcome this caveat we present a novel approach allowing us to derive all HMXB positions. We build the Spectral Energy Distribution (SED) of each HMXB and fit it with a black-body model to compute the distance of each source. Finally, we study this distribution and the correlation with Star Forming Complexes (SFCs) observed in the Galaxy. ",
        "conclusions": "Evaluation of HMXB distribution is now of major interest in order to study in depth the formation of these high energy sources. However, HMXB locations are usually poorly constrained and largely dependent on the determination method. We propose here to determine the location of a whole set of sources using the same approach: a SED fitting of the distance of HMXBs. This method, based on a least-square minimization, enables us to reveal a consistent picture of the HMXB distribution, showing them to follow the spiral arm structure of the Galaxy. The consideration of uncertainties leads to a small error on the source locations and allows us to tackle the study of the correlation with SFC distribution. This study shows that HMXBs are clustered with SFCs and enables for the first time to quantitatively define the cluster size (0.3 kpc) and the distance between clusters (1.7 kpc). The challenge is now to quantify this correlation by taking into account the offset between current spiral density wave position and HMXB positions expected due to the fact that the matter rotation velocity is different to the spiral arm rotation speed."
    },
    "1107/1107.5306_arXiv.txt": {
        "abstract": "We explore the properties of dust and associated molecular gas in 352 nearby ($0.01<z<0.07$) early-type galaxies (ETGs) with prominent dust lanes, drawn from the Sloan Digital Sky Survey (SDSS). Two-thirds of these `dusty ETGs' (D-ETGs) are morphologically disturbed, which suggests a merger origin, making these galaxies ideal test beds for studying the merger process at low redshift. The D-ETGs preferentially reside in lower density environments, compared to a control sample drawn from the general ETG population. Around 80 per cent of D-ETGs inhabit the field (compared to 60 per cent of the control ETGs) and less than 2 per cent inhabit clusters (compared to 10 per cent of the control ETGs). Compared to their control-sample counterparts, D-ETGs exhibit bluer ultraviolet-optical colours (indicating enhanced levels of star formation) and an active galactic nucleus fraction that is more than an order of magnitude greater (indicating a strikingly higher incidence of nuclear activity). The mass of clumpy dust residing in large-scale dust features is estimated, using the SDSS $r$-band images, to be in the range $10^{4.5}$--$10^{6.5}$M$_{\\odot}$. A comparison to the total (clumpy + diffuse) dust masses -- calculated using the far-infrared fluxes of 15 per cent of the D-ETGs that are detected by the \\emph{Infrared Astronomical Satellite (IRAS)} -- indicates that only 20 per cent of the dust is typically contained in these large-scale dust features. The dust masses are several times larger than the maximum value expected from stellar mass loss, ruling out an internal origin. The dust content shows no correlation with the blue luminosity, indicating that it is not related to a galactic scale cooling flow. Furthermore, no correlation is found with the age of the recent starburst, suggesting that the dust is accreted directly in the merger rather than being produced in situ by the triggered star formation. Using molecular gas-to-dust ratios of ETGs in the literature, we estimate that the median current molecular gas fraction in the \\emph{IRAS}-detected ETGs is $\\sim$1.3 per cent. Adopting reasonable values for gas depletion time-scales and starburst ages, the median \\emph{initial} gas fraction in these D-ETGs is $\\sim$4 per cent. Recent work has suggested that the merger activity in nearby ETGs largely involves minor mergers (dry ETG + gas-rich dwarf), with mass ratios between 1:10 and 1:4. If the \\emph{IRAS}-detected D-ETGs have formed via this channel, then the original gas fractions of the accreted satellites are between 20 and 44 per cent. ",
        "introduction": "As endpoints of the hierarchical build-up of mass, early-type galaxies (ETGs) possess a detailed fossil record of their progenitors. The stellar populations of ETGs encode the mass assembly history of galaxies over the lifetime of the Universe, and it is important that we gain a comprehensive understanding of their formation and evolution. While their red optical colours \\citep*[e.g.][]{bow92,ber03,bel04,fab07}, high alpha-element enhancements \\citep*[e.g.][]{tho99,tra00a,tra00b} and obedience of a tight Fundamental Plane \\citep*[e.g.][]{jor96,van96,sag97,for98} indicate that the bulk of their stellar populations ($>80$ per cent) are old, the large scatter in the star-formation-sensitive ultraviolet (UV) colours of ETGs since $z\\sim 1$ is interpreted as evidence for continuous low-level star formation at late epochs \\citep*[][see also \\citealt{fuk04,yi05,jeo07,sal10,cro11}]{kav07,kav08,mar09}. A coincidence between blue UV colours and morphological disturbances indicates that this star formation is merger driven \\citep{sch90,sch92,kav10,kav11}. Furthermore, the paucity of major mergers at late epochs \\citep[see e.g.][]{ste08,jog99,lop10} strongly suggests that the star formation is driven by \\emph{minor} mergers \\citep{kav09,kav11}. Our current understanding of ETG evolution therefore indicates that their underlying stellar populations are old, having been rapidly built at high redshift ($z>1$), while the evolution at late epochs ($z<1$) is dominated by repeated minor-merger events that contribute $<20$ per cent of their stellar mass at the present day.  Notwithstanding the classical notion of ETGs being dry, passively evolving systems, an extensive literature has developed over the past few decades on the interstellar medium (ISM) in these systems. The dust and gas contents of \\emph{very} nearby ETGs have been studied by several authors \\citep*[e.g.][]{tub80,haw81,sad85,ber87,ebn88,kna89,gou95,van95,kna96,fab97,tom00,tra01,com07,cal08,you11}. While the galaxy samples and methodologies in these studies are varied, it is clear that the majority of ETGs in the very nearby Universe show evidence of dust \\citep[e.g.][]{kna89,van95}, the dust masses being generally inconsistent with scenarios in which the dust is supplied purely by stellar mass loss \\citep[e.g.][]{mer98}. Coupled with the frequently observed morphological disturbances and kinematical misalignments between gas and stars in these systems \\citep*[e.g.][]{zei90,sag93,ann10,dav11}, it seems likely that a significant fraction of the interstellar matter has an external origin.  The accumulating evidence for widespread minor-merger-driven star formation \\citep[e.g.][]{kav09,kav11} makes the study of interstellar matter in ETGs all the more compelling. While the properties of the star formation have been quantified with a reasonable degree of precision, less is known about the fuel that drives this star formation. A study of the ISM of ETGs in modern observational surveys is therefore very desirable. In this paper we present a study of dusty ETGs (D-ETGs), drawn from Data Release 7 (DR7) of the Sloan Digital Sky Survey \\citep[SDSS;][]{aba09}, that exhibit prominent, extended dust features in their optical images. In addition to exploring the properties of the dust, the homogeneous nature of the spectrophotometric data from the SDSS allows us to systematically study the star formation, active galactic nucleus (AGN) activity and local environment of the D-ETGs, compare these systems to a statistically meaningful `control' sample drawn from the general ETG population and put some of the results in the literature on a firmer statistical footing. We are also able to explore the properties of the gas in the minor mergers that are responsible for star formation in nearby ETGs, and provide a more complete picture of the star formation activity in massive ETGs at late epochs.  The plan for this paper is as follows. Section~\\ref{section:data} describes the selection of D-ETGs from the SDSS and the construction of a control sample to which their properties are compared. In Section~\\ref{section:morphologies} we discuss the high incidence of disturbed morphologies observed in the D-ETG sample. In Section~\\ref{section:environment} we explore the local environment of these galaxies, and in Section~\\ref{section:sf_agn} we employ UV-optical spectrophotometry to study their star formation and AGN activity. We explore the properties and origin of the dust in Section~\\ref{section:dust} and conclude by summarising our findings in Section~\\ref{section:summary}. ",
        "conclusions": "\\label{section:summary} We have explored the properties of dust and associated molecular gas in ETGs, by studying 352 nearby ($0.01<z<0.07$) ETGs with prominent dust features in the SDSS DR7. The galaxies were identified through direct visual inspection of SDSS images using the GZ2 project. Our analysis has focused on comparing the local environment, star formation and nuclear activity of these D-ETGs to those of a control sample drawn from the general ETG population. We have then explored the properties and origin of the dust and associated molecular gas in these systems, and studied the processes that lead to their formation.  Two-thirds of the D-ETGs show morphological disturbances at the depth of the standard SDSS images, suggesting a merger origin for these galaxies. D-ETGs reside preferentially in low-density environments, with 80 per cent inhabiting the field (compared to 60 per cent of their control-sample counterparts) and less than 2 per cent inhabiting clusters (compared to 10 per cent of their control-sample counterparts). D-ETGs show bluer UV-optical colours compared to the control ETG population, indicating the presence of enhanced levels of star formation activity. An analysis using optical emission-line ratios indicates that the fraction of objects classified as star forming and composite are, respectively, factors of 3 and 8 higher in the D-ETGs compared to their control-sample counterparts. The fraction of Seyferts is more than an order of magnitude higher, while the fraction of galaxies classified as quiescent is correspondingly an order of magnitude smaller in the D-ETG population compared to that in the control sample. The presence of prominent dust features in an ETG therefore significantly increases its chances of hosting both star formation and nuclear activity, with nuclear activity being the dominant gas ionization mechanism.  We have explored the dust content of our D-ETGs, both in terms of the clumpy dust mass contained in the large-scale dust features (visible in the optical SDSS images) and in terms of the diffuse dust that may exist in the ISM. The clumpy dust masses span the range $10^{4.5}$--$10^{6.5}$M$_{\\odot}$ and are typically several factors higher than the maximum dust contribution expected from stellar mass loss from the old, underlying stellar populations that dominate ETGs. The clumpy dust mass does not correlate with the blue luminosity of the galaxies, indicating that galactic-scale cooling flows are unlikely to be responsible for the dust content. Similarly, no correlation is found between the clumpy dust content and the age of the recent starburst, suggesting that the dust is not produced in situ from the star formation triggered by the merger event, but is rather accreted directly during the merger.  The clumpy dust masses are strictly lower limits to the total dust masses because they do not include diffuse dust which may be intermixed with the stars and is essentially invisible in the optical images. In the 15 per cent of our D-ETGs that are detected by \\emph{IRAS}, we have estimated the `total' (clumpy + diffuse) dust mass using the \\emph{IRAS} 60- and 100-$\\mu$m fluxes. The average discrepancy between the clumpy and \\emph{IRAS}-derived dust mass in these D-ETGs is a factor of 4.5. In other words, only $\\sim$20 per cent of the dust mass in these objects is contained in the large-scale dust features that are observed in the optical images. Nevertheless, the \\emph{IRAS} dust masses show the same lack of correlation with galaxy properties and starburst ages as their clumpy counterparts.  The correspondence between dust and cold gas allows us to explore the properties of the cold gas reservoirs in our D-ETGs. Restricting ourselves to the \\emph{IRAS}-detected galaxies (in which we have a reasonable estimates of the total dust mass), we employ the average molecular G/D of ETGs in the literature ($\\sim 1190$) to estimate the current molecular gas fractions of these objects. The median value of the current gas fraction is $\\sim$ 1.3 per cent. Given the current molecular gas fraction, an estimate of the age of the starburst and a characteristic gas depletion (star formation) time-scale, we have estimated the \\emph{initial} molecular gas fraction in each \\emph{IRAS}-detected D-ETG, with a median value of $\\sim$4 per cent. Given recent work that suggests that the merger activity in nearby ETGs largely involves minor mergers with mass ratios between 1:10 and 1:4 \\citep[e.g.][]{kav09}, a typical initial molecular gas fraction of 4 per cent suggests that the gas-rich satellites accreted by the \\emph{IRAS}-detected ETGs are likely to have had molecular gas fractions between $\\sim$20 and 44 per cent (assuming that the ETGs themselves are gas and dust free).  While recent UV-optical studies have measured the continuous (minor-merger-driven) star formation in ETGs over the last 8 billion years, less is known about the dust and associated gas that drives this star formation. This paper has offered insights into the ISM of nearby ETGs that are merger remnants and, in particular, the initial gas masses plausibly brought in by the satellites that drive the star formation. The galaxy sample studied here provides an ideal basis for future studies of D-ETGs using new instruments such as \\emph{Herschel} and the Atacama Large Millimeter Array (ALMA) to probe the properties of the ISM of nearby ETGs and achieve a fuller understanding of the gas that drives the star formation in these galaxies at late epochs."
    },
    "1107/1107.5130_arXiv.txt": {
        "abstract": "We consider a relativistically moving blob consisting of an isotropic electron distribution that Compton-scatters photons from an external isotropic radiation field. We compute the resulting beaming pattern, i.e. the distribution of the scattered photons, in the blob frame as well as in the observer's frame by using the full Klein-Nishina cross section and the exact incident photon distribution. In the Thomson regime the comparison of our approach with \\citet{Dermer1995} results in concurrent characteristics but different absolute number of the scattered photons by a factor of $f_{corr}=3.09$. Additionally, our calculation yields a slightly lower boost factor which varies the more from the corresponding value in \\citet{Dermer1995} the higher the spectral index $p$ of the electron distribution gets. ",
        "introduction": "Blazars are flat radio spectrum active galactic nuclei that are dominated by the nonthermal emissions of superluminal moving bulk plasma blobs emitted from the core \\citep{Zensus1987}. Their strong high energy emissions show a characteristic double bump structure in their spectra \\citep{Sanders1989}. In leptonic models, the lower energy bump is caused by Synchrotron emission while the bump at higher energies is produced by inverse Compton-scattering. According to Synchrotron Self Compton models (SSC) Synchrotron photons are produced by accelerated particles in the jet, while gamma ray photons arise from the Inverse Compton-scattered Synchrotron photons \\citep{Maraschi1992,Marscher1994}. However, external photons that are not produced inside the jet can also be scattered by the Inverse Compton effect and can contribute to the emission. External Compton models (EC) consider photons from the accretion disk or the Broad Line Region \\citep{Blandford1993,Dermer1992,Dermer1993,Sikora1994} while conventional one zone SSC models consider the emission in the blob as homogeneous. These models assume that the variability of blazar emissions is limited by the size of the blob. In contrast, two zone SSC models \\citep{Weidinger2010} allow for shorter variabilities but only spatially resolved models are suitable for even shorter variabilities. They take angular dependent properties of the jet into account like the direction of magnetic fields and spatial variations of electron and photon densities inside the blob.  In this paper we consider spatially resolved SSC models with an EC component and are interested in ascertaining the angular radiation characteristics of the EC component. \\cite{Dermer1995} assumed that this beaming pattern is produced by a relativistic moving blob consisting of isotropic distributed electrons that Thomson-scatter photons from an external isotropic radiation field. As a result the isotropic external photons are more strongly boosted than isotropically distributed photons in a blob within the jet. While Dermer's calculations are based on ultrarelativistic electrons and scattering in the Thomson regime, in this paper we will consider all possible electron energies obeying the Thomson as well as the Klein-Nishina regime. Using the exact Klein-Nishina cross section we compute the distribution of scattered photons numerically, examine the dependencies of physical parameters and determine the resulting intensity for an electron distribution given by a power law. Such a calculation has already been executed by \\cite{Georganopoulos2001} but in contrast to them we do not use the formula of the  scattered photons by \\cite{Jones1968} that averages over the photons' emergent angles. In fact we will inherit the first steps of Jones exact calculation, i.e. not relying on the head on approximation, and continue using exact expressions.  We organize the remaining paper as follows. In Section \\ref{sec_calc} the necessary steps are shown to calculate the scattered spectrum exactly. In Section \\ref{sec_results} the results in the blob frame as well in the observer's frame are presented and compared to both \\cite{Dermer1995} and \\cite{Georganopoulos2001}. ",
        "conclusions": "\\label{sec_results} The characteristic spectrum of the Compton-scattered photons is a broadened peak if we assume an isotropic electron distribution in the blob frame. When plotting the spectrum on a double logarithmic scale graph, two characteristic areas show up. These are the regime of elastic scattering, i.e. the Thomson regime, and the regime of inelastic scattering, i.e. the Klein-Nishina regime. Defining the electron distribution by a spectral index $p$ the differential number of photons in the Thomson regime depends on $\\alpha^{-(p-1)/2}$. In the Klein-Nishina regime the spectral index is characterized by a higher spectral index.  \\subsection{Blob frame} Choosing an electron distribution of $n_e(\\gamma)=\\gamma^{-2.2}$, the plot of the differential number of scattered photons displays the transition from the Thomson (elastic scattering) to the Klein-Nishina regime (inelastic scattering) that is displayed as a sharp break in Figure \\ref{fig_comp_dermer}. The upper limit of the Thomson regime depends on the initial energy in the observer's frame because depending on the incident angle of the external photons they loose or win energy due to the change of the frame of reference. The maximal electron energy $\\gamma_{T,max}$ that allows scattering in the Thomson regime determines this upper limit: $\\alpha_{1,min} \\gamma_{T,max} <1$. Hence, the maximal photon energy that results from Thomson scattering is defined by the maximal electron energy, i.e. $\\alpha_{T,max}=4 \\alpha_{1,min} \\gamma_{T,max}^2$. The higher $\\Gamma$, the more the incident photon distribution is peaked, the higher the maximum energy $\\alpha_{1,max}=2\\Gamma \\alpha_1^*$ and the lower $\\gamma_{T,max}$ and the Thomson limit.  \\subsection{Observer's frame} \\begin{figure}[htb] \\centering \\includegraphics[width=11cm]{Fig2.pdf} \\caption{Comparison of the gamma ray emissivities assuming different maximal electron energies $\\gamma_{max}=10^6$, $10^8$, $10^{10}$, $10^{12}$. Emissions are presented in the observer frame assuming a spectral index of the electron distribution of $p=2.2$, $n_e^0=1$, $n_{ph}=1$ and an incident photon energy of $\\alpha^*_1=10^{-9}$. Emissions are calculated in the observer frame and correspond to a viewing angle of $0^{\\circ}$ if $\\Gamma=20$.} \\label{fig_energies} \\end{figure} \\begin{figure}[htb] \\centering \\includegraphics[width=11cm]{Fig3.pdf} \\caption{Comparison of the gamma ray emissivities using angle resolved and angle averaged calculations. Emissions are presented in the observer frame assuming a spectral index of the electron distribution of $p=2.2$, $n_e^0=1$, $n_{ph}=1$ and an incident photon energy of $\\alpha^*_1=10^{-9}$. Emissions are calculated in the observer frame and correspond to a viewing angle of $0^{\\circ}$ if $\\Gamma=20$.} \\label{fig_comp_georg} \\end{figure} Boosting the resulting spectrum into the observer's frame enhances the differential number of photons emerging at $\\theta^*_{\\alpha}\\rightarrow \\pi$ and decreases those emerging at $\\theta^*_{\\alpha}\\rightarrow 0$. The main characteristics of the spectrum - in particular the break - remain. In Figure \\ref{fig_energies} the development of the break is shown by choosing different maximal electron energies of $\\gamma_{max}=10^{6}$, $10^8$, $10^{10}$, $10^{12}$. The higher the electron energy, the higher the maximal energy of the scattered photons. In the Klein-Nishina regime photons are scattered to energies up to $\\alpha^*_{KN,max}=2\\Gamma\\gamma_{max}$ while in the Thomson regime up to $\\alpha^*_{T,max}=2\\Gamma\\ 4 \\alpha_{1,min} \\gamma_{T,max}^2$. The break gradually appears as scattering is dominated by the Klein-Nishina regime. It turns out that this break is sharper than in \\citet{Georganopoulos2001}. As we do not use any approximations (e.g. head on approximation, outgoing photons are directed along the direction of the scattering electrons, i.e. $\\Omega_{\\alpha}=\\Omega_e$) when calculating the spectrum the sharper break is due to the angle averaged solution of \\citet{Georganopoulos2001} (introduced by \\citet{Jones1968}). By taking the average of all emergent photon angles in the blob frame we studied how angle averaging effects the shape of the break (cf. Figure \\ref{fig_comp_georg}). The break becomes less sharp and is similar to the shape of the break in \\citet{Georganopoulos2001}. The number of photons of the angle averaged solution is lower than those of the exact calculation due to the decreasing photon number at lower emerging angle $\\theta^*_{\\alpha}$.  \\subsection{Comparison with \\citet{Dermer1995}}  \\begin{figure}[htb] \\centering \\includegraphics[width=11cm]{Fig4.pdf} \\caption{Comparison between the gamma ray emissivities using the full Klein-Nishina cross section and Dermer's approximation. Emissions are calculated in the blob frame and correspond to a viewing angle of $0^{\\circ}$ if $\\Gamma=20$. The spectral index for the electron distribution runs $p=2.2$, $n_e^0=1$, $n_{ph}=1$, $1\\leq \\gamma\\leq 10^{12}$.} \\label{fig_comp_dermer} \\end{figure} In the following section we compare our results to those of  \\citet{Dermer1995} with respect to the regime of elastic scattering (Thomson regime). We observe the differential scattered photon number to be larger than predicted by \\citet{Dermer1995} (cf. Figure \\ref{fig_comp_dermer}). Dermer considers only the emergent photons that are scattered to the maximal energy $\\alpha_{max}=4\\alpha_1 \\gamma^2$. This fact explains the observed difference since in our work all photons with energies ranging from the initial photon energy to the maximum energy $\\alpha_{max}$ are also considered. We obtain Dermer's result if we only take into account the photons that are scattered to the maximum energy. The distribution of scattered photons in the observer's frame as well as the difference to Dermer highly depends on the observing angle and the velocity of the blob. Depending on the emergent angle $\\theta_{\\alpha}$ and the gamma factor of the blob $\\Gamma$ the relative difference $\\epsilon$ of the differential number of scattered photons between our exact calculation and \\citet{Dermer1995} is shown in Figure \\ref{fig_valid_dermer} as a contour plot. The mean value for all relative differences for $2.00\\leq\\epsilon\\leq2.12$ is given by $\\epsilon_{lim}=2.09$ and hence the correcting factor runs $f_{corr}=1+\\epsilon_{lim}=3.09$. The limits for the gamma factor of the blob $\\Gamma$ and the emergent angle $\\theta_{\\alpha}$ can be deduced from the values in Figure \\ref{fig_valid_dermer}. \\begin{figure}[htb] \\centering \\includegraphics[width=12cm]{Fig5.pdf} \\caption{Relative difference $\\epsilon$ of the differential number of scattered photons between this exact calculation and \\citet{Dermer1995}, dependent on the emergent angle $\\theta$ and the gamma factor of the blob $\\Gamma$. The lines are contour lines and show together with the coulors the value of the relative difference.} \\label{fig_valid_dermer} \\end{figure}  We find that for higher emergent angles $\\theta_{\\alpha}>2.09\\ \\mathrm{rad}$ and higher blob velocities $\\Gamma>2.8$ our differential number of scattered photons is higher by a factor of $f_{corr}=3.09$ than calculated in \\citet{Dermer1995}. Hence, Dermer's approximations can be used for high Doppler factors and small observing angles, i.e. $\\theta_{\\alpha}=\\pi$. But for higher observing angles and smaller Doppler factors the deviation increases and the Klein-Nishina cross section and the exact incident photon distribution must be considered. Considering the boost factor that is needed to transform the differential number of scattered photons from the blob frame into the observer's frame, we computed a lower  boost factor than derived in \\citet{Dermer1995}. Analogously to \\citet{Dermer1995} the boosting depends on the spectral index $p$ of the electron distribution. If $a_{\\alpha}$ and $a_{\\theta_{\\alpha}}$ describe the exponents of the Dopplerfactor of the respective quantities, the found boost factor goes as $D^{3+a_{\\alpha}+1+a_{\\theta_{\\alpha}}}$. $a_{\\alpha}$ coincidences with Dermer's $a=(p-1)/2$ and $a_{\\theta_{\\alpha}}$ reaches a smaller value. This difference can be explained by Dermer's use of the head on approximation. Dermer uses this approximation for all scattering processes but actually it is not valid for low energetic electrons. Furthermore the number of low energetic electrons coincides with the decrease of the deviation when smaller spectral indices of the electron distribution are considered.  \\subsection{Conclusions} Comparing our work to \\citet{Georganopoulos2001} we did not use the head on approximation as well as the calculated scattered spectrum by \\citet{Jones1968} that averages over all angles of the emergent photons. In \\citet{Georganopoulos2001} the only angle dependent quantity is the Doppler factor while we computed directly the angle dependent scattered spectrum. While the rough shape and peak energies coincidences with \\citet{Georganopoulos2001} we find as already mentioned a smaller boost factor than \\citet{Dermer1995} and \\citet{Georganopoulos2001}. Due to the exact calculation the break in our spectrum is much sharper than in \\citet{Georganopoulos2001}. But taking the average of all angles of the emergent photons softens the break and approaches the results of \\citet{Georganopoulos2001}.  Summing up, in contrast to \\citet{Dermer1995} we observe a higher differential number of scattered photons in the regime of elastic scattering. Furthermore a lower boost factor was computed because we considered the exact distribution of the incident photon field. Despite this difference Dermer's approximation is applicable for high Doppler factors and small observing angles. But for high inclination angles (e.g. M87, \\citealt{Giovannini2010}) or small Doppler factors (e.g. BL Lacertae 3C371.0 or RG 3C84, \\citealt{Hovatta2009}) the calculations of this work should be used. Furthermore for high electronic spectral indices Dermer's boost factor overvalues and a too large number of scattered photons is yielded."
    },
    "1107/1107.0009_arXiv.txt": {
        "abstract": "We compare models for Type Ia supernova (\\snia) light curves and spectra with an extensive set of observations. The models come from a recent survey of 44 two-dimensional delayed-detonation models computed by \\cite{KRW09}, each viewed from multiple directions. The data include optical light curves of 251 \\sneia, some of which have near-infrared observations, and 2231 low-dispersion spectra from the Center for Astrophysics, plus data from the literature. These allow us to compare a wide range of \\snia\\ models with observations for a wide range of luminosities and decline rates. The analysis uses standard techniques employed by observers, including MLCS2k2, SALT2, and SNooPy for light-curve analysis, and the Supernova Identification (SNID) code of Blondin \\& Tonry for spectroscopic comparisons to assess how well the models match the data. The ability to use the tools developed for observational data directly on the models marks a significant step forward in the realism of the models. We show that the models that match observed spectra best lie systematically on the observed width-luminosity relation. Conversely, we reject six models with highly asymmetric ignition conditions and a large amount ($\\gtrsim1$\\,M$_{\\sun}$) of synthesized \\nifs\\ that yield poor matches to observed \\snia\\ spectra. More subtle features of the comparison include the general difficulty of the models to match the $U$-band flux at early times, caused by a hot ionized ejecta that affect the subsequent redistribution of flux at longer wavelengths. The models have systematically higher velocities than the observed spectra at maximum light, as inferred from the Si\\two\\,\\l6355 line. We examine ways in which the asymptotic kinetic energy of the explosion affects both the predicted velocity and velocity gradient in the Si\\two\\ and Ca\\two\\ lines. Models with an asymmetric distribution of \\nifs\\ are found to result in a larger variation of photometric and spectroscopic properties with viewing angle, regardless of the initial ignition setup. We discuss more generally whether highly anisotropic ignition conditions are ruled out by observations, and how detailed comparisons between models and observations involving both light curves and spectra can lead to a better understanding of \\snia\\ explosion mechanisms. ",
        "introduction": "\\label{sect:intro}  Type Ia supernovae (\\sneia) play a major role in many astrophysical phenomena. They produce a large fraction of iron in the universe (e.g., \\citealt{Truran/Cameron:1971}), heat the interstellar medium (e.g., \\citealt{Ciotti/etal:1991}), and form an endpoint of binary star evolution (e.g., \\citealt{Iben/Tutukov:1984}). \\sneia\\ provide the most reliable and precise cosmological distances to establish the acceleration of cosmic expansion \\citep{R98,P99}.  Despite their astrophysical importance, however, they remain enigmatic objects. There is a general consensus that they result from the thermonuclear disruption of a carbon-oxygen white dwarf (WD) star \\citep{Hoyle/Fowler:1960} approaching the Chandrasekhar mass ($M_{\\rm Ch}\\approx1.4$\\,M$_{\\sun}$), either through accretion from a non-degenerate binary companion (the ``single-degenerate'' scenario), or through merger with another WD (the ``double-degenerate'' scenario; \\citealt{Iben/Tutukov:1984,Webbink:1984}). Which of these two possibilities constitutes the dominant (or sole) progenitor channel for \\sneia\\ is still debated (see \\citealt{Howell:2010} for a recent review).  The explosion mechanism itself is also largely unknown (see \\citealt{Hillebrandt/Niemeyer:2000} for a review). In the preferred ``delayed-detonation'' model \\citep{Khokhlov:1991}, the burning starts as a turbulent subsonic deflagration near the WD center and transitions to a supersonic detonation near its surface. The deflagration pre-expands the WD so that the subsequent detonation does not burn the entire star to nuclear statistical equilibrium (NSE) material (including \\nifs\\ to power the light curve), but instead synthesizes appropriate fractions of high-velocity ($\\sim10000$\\,\\kms) intermediate-mass elements (IME; such as Mg, Si, S, Ca, etc.) needed to reproduce the observed spectra. The first simulations were carried out in 1D, but recent studies show that multi-dimensional simulations are needed to capture hydrodynamical instabilities (e.g., \\citealt{Gamezo/Khokhlov/Oran:2005,Roepke/Niemeyer:2007}) and to provide a physical basis for the transition from deflagration to detonation. (e.g., \\citealt{Woosley:2007,Roepke:2007,Woosley/etal:2009}).  The empirical relation between the peak luminosity and the width of the light-curve (the so-called width-luminosity relation, or WLR; \\citealt{Pskovskii:1977,Phillips:1993}), instrumental to the use of \\sneia\\ as distance indicators, can be physically interpreted in terms of (1) varying opacity with the amount of synthesized \\nifs\\ \\citep{Hoeflich/Khokhlov:1996}, (2) varying mass of the progenitor WD \\citep{Pinto/Eastman:2000a}, or (3) the {\\sc iii}$\\rightarrow${\\sc ii} (i.e. doubly- to singly-ionized) recombination timescale of iron-group elements in the SN atmosphere \\citep{Kasen/Woosley:2007}. More detailed calculations are also needed to explain the observed scatter in the WLR.  In a recent paper, \\cite{KRW09} (hereafter KRW09) conducted a 2D survey of delayed-detonation models, in which they varied the radial/angular distribution and number of ignition points in a Chandrasekhar-mass C+O WD star, as well as the criterion for deflagration-to-detonation transition. By synthesizing light curves and spectra for different viewing angles, they were able to broadly reproduce both the observed width-luminosity relation (for all but subluminous \\sneia) and its scatter, illustrating the importance of multi-dimensional computations to reproduce observed trends in \\snia\\ properties. Furthermore, they showed that variations in the metallicity of the progenitor WD affect both the slope and normalization of the WLR, and that ignoring these effects could potentially lead to systematic overestimates of $\\sim2$\\%\\ on distance determinations to \\sneia.  To study such subtle effects, one must ensure the models reproduce all observed trends in some detail, and not just the width-luminosity relation. This is precisely what we set out to do in the present study, where we conduct an extensive analysis of the delayed-detonation models of KRW09 through a detailed and direct comparison with observations of \\sneia. We present the models and data in \\S~\\ref{sect:modeldata}, and our methods for quantitatively evaluating each model in \\S~\\ref{sect:method}. We then proceed to a detailed comparison of their photometric (\\S~\\ref{sect:phot}) and spectroscopic (\\S~\\ref{sect:spec}) properties with actual data. We discuss whether explosion models with asymmetric ignition conditions are ruled out by observations in \\S~\\ref{sect:disc}, and conclude in \\S~\\ref{sect:ccl}. ",
        "conclusions": "\\label{sect:ccl}  We have presented a detailed comparison of a recent survey of 2D delayed-detonation explosion models by KRW09 with observations of Type Ia supernovae. We apply standard methods used by \\snia\\ observers to compare the model light curves and spectra with empirical templates. This represents a significant step forward in the realism of the models. Running several light-curve fitters (MLCS2k2, SALT2, SNooPy) on synthetic $(U)BVRI$ light curves, we find some tension between the light-curve shape of the models and actual data, the models having longer rise times. Based on cross-correlations with a library of \\snia\\ spectra, we quantified the overall resemblance of individual models to observed \\sneia, and found that the best models/viewing angles lied systematically on the observed width-luminosity relation.  Comparison of several photometric properties of the models (rise times, maximum-light colours and their evolution with time) shows a broad agreement with observations, but reveals some problems with flux redistribution from the near-UV to the near-IR bands, a key mechanism needed to explain both the width-luminosity relation and the secondary maxima in the NIR light curves, and mediated by the {\\sc iii}$\\rightarrow${\\sc ii} recombination timescale of iron-group elements \\citep{Kasen:2006,Kasen/Woosley:2007}. Subsequent investigation of spectra for a subset of selected models confirmed the excess of $U$-band flux in the models at early times, likely caused by a hot ionized ejecta and subsequent lack of absorption by Fe\\two/Co\\two. Interestingly, one of our selected models (DD2D\\_iso\\_08\\_dc3) that shows the best overall agreement in optical and NIR colour evolution with observations lies off the width-luminosity relation (i.e. the colours match, but not the luminosity). This reveals one limitation of our approach, which relies on maximum-light spectra to rank the different models, whereas a combination of photometric and spectroscopic properties is needed for a proper evaluation.  Comparison of maximum-light spectra show the models have systematically large absorption velocities (most visible in the Si\\two\\,\\l6355 line), affecting the relative shapes and strengths of spectral features and smoothing out small substructures observed in iron-dominated absorption complexes at $\\sim4300$\\,\\AA\\ and $\\sim4800$\\,\\AA. Consequently, correlations between several spectroscopic indicators and \\dmft\\ decline rate have a much larger scatter in the models. The relation found by \\cite{Foley/Kasen:2011} between Si\\two\\,\\l6355 absorption velocity and intrinsic $B-V$ colour (redder \\sneia\\ having larger $|v_{\\rm abs}|$) is also not reproduced in the models, most showing a weak correlation in the opposite direction. However, the observed correlation is weak and subject to a large uncertainty given the errors on intrinsic $B-V$ colour inferred from the data. Nonetheless, we identify a trend of larger absorption blueshifts for higher kinetic energy for models in which only the criterion for deflagration-to-detonation transition is varied.  The overall evolution of the model spectra compares well with observations, as illustrated by the comparison of synthetic spectra for model DD2D\\_iso\\_06\\_dc2 between $-10$\\,d and +20\\,d from $B$-band maximum with observed spectra of SN~2003du, but several discrepancies characteristic of most models are apparent. The synthetic spectra are too blue (i.e. too hot) at early times, and the ionization that ensues affects their subsequent evolution. Most notably, the models appear to lack an absorption feature around $\\sim5000$\\,\\AA\\ (attributed to Fe\\two/Co\\two), and fail to reproduce the strong emission feature at $\\sim5800$\\,\\AA\\ (attributed to Na\\one\\,D) from +10\\,d onwards. Non-LTE calculations of \\snia\\ spectra also fail to reproduce this line \\citep{Baron/etal:2006}.  The evolution of the Ca\\two\\,\\l3945 absorption velocity with time exhibits a strong diversity in the models which contrasts with the steady and smooth decrease seen in the data. Model DD2D\\_asym\\_01\\_dc3 has an asymmetric \\nifs\\ distribution and shows all types of behaviour depending on the viewing angle, due to the varying radial Ca distribution with different lines of sight. The evolution of the Si\\two\\,\\l6355 velocity shows little variation before maximum light, while the post-maximum evolution (the velocity gradient) appears conditioned by the kinetic energy of the explosion, affecting the abundance and radial distribution of silicon. While this does not contradict the recent findings of \\cite{Maeda/etal:2010c}, who associate the observed diversity in velocity gradients with viewing angle effects in off-center explosions, it shows that the interpretation of such gradients depends on more than a single parameter of the explosion.  We reject six models of KRW09 with highly asymmetric ignition conditions and are characterized by large amounts ($\\gtrsim 1$\\,\\msun) of \\nifs. We do not reject off-center delayed-detonation models for \\sneia\\ as a whole, but note the extreme sensitivity of the amount and distribution of burning products in the deflagration phase to the initial distribution of ignition points \\citep[see also][]{Livne/Asida/Hoeflich:2005}.  Throughout this paper we have focused on discrepancies between the models and observations more than we have highlighted their mutual agreement, but this merely results from the unprecedented level of detail of our study. Such detail is necessary to use the predictive power of the models to provide a physical basis to some observed trends, as well as use the data to impose meaningful constraints on the models. The 2D delayed-detonation models of KRW09 have a degree of fidelity which makes them amenable to the same analysis we use on observations of \\sneia, but they still require some adjustments to accurately match the data. The ability to reproduce the bolometric/multi-band light curves and the width-luminosity relation is a necessary but not a sufficient condition for a model to be considered a valid approximation of real \\sneia. Further insights from three-dimensional hydrodynamical simulations, more accurate nucleosynthetic post-processing, and full non-LTE radiative transfer calculations are all part of the solution. We are confident that a detailed comparison of light curves and spectra from grids of models using the framework developed in this paper will lead to a better understanding of \\snia\\ explosion mechanisms."
    },
    "1107/1107.5901_arXiv.txt": {
        "abstract": "A comprehensive search for variable and transient radio sources has been conducted using $\\sim$~55,000 snapshot images of the {\\sl FIRST} survey. We present an analysis leading to the discovery of 1,627 variable and transient objects down to mJy levels over a wide range of timescales (few minutes to years). Variations observed range from 20\\% to a factor of 25. Multi-wavelength matching for counterparts reveals the diverse classes of objects exhibiting variability, ranging from nearby stars and pulsars to galaxies and distant quasars. Interestingly, more than half of the objects in the sample have either no classified counterparts or no corresponding sources at any other wavelength and require multi-wavelength follow-up observations. We discuss these classes of variables and speculate on the identity of objects that lack multi-wavelength counterparts. ",
        "introduction": "The first celestial radio source detected -- the Sun -- was discovered as a consequence of its variability \\citep{hey46}. Nonetheless, while variability has continued to be a source of discovery and insight in the radio regime, few dedicated searches for variables and transients have been undertaken, primarily because the observing times required to conduct deep wide-field surveys are large and current instruments tend to have poor figures of merit for variability studies \\citep{cor04}.  Indeed, one of the key scientific aims of the next generation of radio telescopes such as LOFAR, ASKAP, MWA and, eventually, the SKA is a study of radio variability. As discussed by \\citet{hes09}, LOFAR will survey the sky for pulsars and fast transients at low frequencies (30-240~MHz) with an emphasis on timescales of less than a second. CVs, X-ray binaries, GRBs, SNe, AGN, flare stars, exoplanets and many more new phenomena, as well as extrinsic causes such as interstellar and interplanetary scintillation, are expected to contribute to its inventory of radio transients and variables \\citep{fen08}. The Australian SKA Pathfinder \\citep[ASKAP;][]{joh07,joh08} is designed to achieve very high survey speeds and noise levels down to $\\sim 10-100$~$\\mu$Jy in an hour in the 1~GHz band, and will shed light on GRBs, radio supernovae (RSNe), Intra-Day Variables (IDVs), etc. The Murchison Widefield Array (MWA) will operate in the 80-300~MHz range carrying out a blind search for variables and transients in addition to targeting explosive events, stellar and planetary phenomena, and compact objects over a range of timescales from nanoseconds to years \\citep{lon09}. The EVLA \\citep{nap06,rup00} will offer vastly improved sensitivity along with dynamic scheduling, providing a host of new capabilities for transient and variable searches. Finally, a search for radio variables and transients also forms an important part of the scientific objectives of the Square Kilometer Array (SKA). The SKA is expected to discover a number of classes of variable radio objects such as pulsars and magnetars, GRBs that are $\\gamma$-ray loud and $\\gamma$-ray quiet in both afterglow and prompt emission, sub-stellar objects such as brown dwarfs and exoplanets, microquasars, and potentially new classes of astrophysical phenomena \\citep{cor04,cor08,wil04}.  Most of the research on radio variability to date has focused on bright radio samples (e.g., \\citet{gre86}, $S>0.4$~Jy at 408~MHz; \\citet{lis01}, $S>0.4$~Jy at 5~GHz; \\citet{all03}, $S>1.3$~Jy at 5~GHz). At these flux densities, the radio source population is dominated by AGN, while at fainter flux densities ($\\lesssim$~1~mJy), it is dominated by star-forming galaxies \\citep{win99,ric99,hop00}. Research on the variability of faint radio sources has so far yielded small samples as a consequence of the small areas searched and the long integration times required. \\citet{car03} observed the Lockman Hole region on timescales of 19 days and 17 months down to 0.1~mJy at 1.4~GHz and discovered nine variable sources, providing an upper limit to the areal density of variable radio sources, as well as constraints on the beaming angle of GRBs and confusion limits relevant to searches for GRB afterglows. Archival VLA calibrator observations spanning 22 years with 944 independent epochs have been used by \\citet{bow07} to search for radio transients over a single field with half-power beam widths  of 9\\farcm0 and 5\\farcm4 at 5~GHz and 8.4~GHz, respectively. The typical observation's flux density threshold is $300$~$\\mu$Jy. They detect eight transients in single epochs and two transients in two-month averages of the data. Two of the transients were identified as RSNe, while the absence of optical counterparts for the remainder offers a wide variety of possibilities including Orphaned GRB Afterglows (OGRBA), stellar sources, propagation effects, microlensing events, or perhaps mechanisms heretofore unknown.  In the Galaxy, single radio bursts detected from the Parkes Multibeam Pulsar Survey revealed a new population of neutron stars: Rotating Radio Anomalous Transients \\citep[RRATs;][]{mcl06}.  Galactic Center Radio Transients \\citep[GCRTs;][]{hym02,hym05,hym09} have diverse light curves with outburst periods and burst durations varying from minutes to months. Lacking counterparts at other wavelengths, their hosts, and the physical mechanisms involved, remain a mystery.  Two large, sensitive, wide-field surveys exist at radio wavelengths. The NRAO VLA Sky Survey \\citep[NVSS;][]{con98} covered the entire sky north of -40\\arcdeg declination at 1.4~GHz  with a resolution of 45\\arcsec and an rms noise of 0.45~mJy. The Faint Images of the Radio Sky at Twenty-cm survey \\citep[{\\sl FIRST};][]{bec95} has covered $\\gtrsim$~9,000~deg$^2$ of the sky at 1.4~GHz with a uniform rms of 0.15~mJy and an angular resolution of 5\\farcs4. A number of authors have used these two surveys in their overlap region to search for radio transients. \\citet{lev02} tried to constrain the detection rates of orphaned radio afterglows often associated with GRBs. \\citet{gal06} attempted to characterize the sample generated by \\citet{lev02}. They report the detection of a radio SN in a nearby galaxy and the detection of a source with no optical counterpart, concluding the latter is unlikely to be associated with a GRB. They also place tighter constraints on the beaming factor of GRBs and a limit on total rate of nearby relativistic explosions, implying that most core collapse SNe do not eject unconfined relativistic outflows. These studies have been complicated by the mismatched resolutions and flux density sensitivities of the NVSS and the {\\sl FIRST} surveys.  \\citet{dev04} have used $\\sim$~120~deg$^2$ of the {\\sl FIRST} survey data near $\\delta=0$\\arcdeg taken in 1995 and then repeated in 2002 to study the optical properties of sources that exhibit significant radio variability at 1.4~GHz over this seven-year interval. They find 123 variable objects with flux densities ranging from $\\sim$~2~-~1000~mJy. They conclude that there is a higher fraction of quasars in the sample of variables compared to the non-varying sample.  More recently, \\citet{ofe10} have raised a very interesting possibility that the mysterious transients of \\citet{bow07} observed out of the plane of the Galaxy, based on their areal number density, duration and energy characteristics, could have a progenitor population consistent with being Galactic isolated old neutron stars with mean distances of the order of kpc. X-ray follow-up \\citep{cro11} of the transients reported by \\citet{bow07} mostly resulted in non-detections and the X-ray flux upper limits imply consistency of the progenitor population with the above possibility besides being consistent with extreme flare stars at $\\gtrsim$~1~kpc or less extreme flares from brown dwarfs at distances of $\\sim$~100~pc.  Here, we use the data from the {\\sl FIRST} survey to create the largest, unbiased sample of variable radio sources to date down to a sensitivity level of a few mJy. In \\S\\ref{FIRST}, we describe the attributes of the {\\sl FIRST} survey relevant to a study of radio variability. In \\S\\ref{dataanalysis}, we present our approach to extracting a sample of variables from the more than two million individual source observations contained in the database. We then describe (\\S\\ref{selection}) our refinement of the list of variable candidates, the parameterization of their variability, and a summary of their properties. In \\S\\ref{catalog}, \\S\\ref{crossid} and \\S\\ref{summary}, we present the results of our analysis, including the cross-identification of our sample of transients and variables with existing data at optical wavelengths (and in other spectral regimes), and a discussion of the classes of variables identified, as well as the new populations which might exist among the unidentified radio variables. ",
        "conclusions": "\\label{summary}  After analyzing $\\sim 55,000$ snapshot images from the {\\sl FIRST} survey of the radio sky covering 8444 deg$^2$ with sensitivity down to 1 mJy for variability and transient phenomena, we have assembled a sample of 1627 sources that are significantly variable. This sample was matched with multi-wavelength catalogs to identify counterparts. We found 13 radio pulsars, 53 SDSS spectroscopic QSOs, 67 SDSS photometric QSOs, 489 galaxies, 5 stars, 123 optically detected but unclassified sources and 877 objects lacking optical counterparts. The unidentified sources mostly occupy the parameter-space of galaxies and QSOs but there are a few notable outliers. Follow-up observations of several of these source classes would likely be fruitful.  Our study has shown that there is much to be discovered in the dynamic radio sky. Exploration by the next-generation instruments sampling different portions of the spectral and temporal domains will prove to be highly productive."
    },
    "1107/1107.1075_arXiv.txt": {
        "abstract": "We quantify the potential for testing MOdified Newtonian Dynamics (MOND) with LISA Pathfinder (LPF), should a saddle point flyby be incorporated into the mission. We forecast the expected signal to noise ratio (SNR) for a variety of instrument noise models and trajectories past the saddle. For standard theoretical parameters the SNR reaches middle to high double figures even with modest assumptions about instrument performance and saddle approach. Obvious concerns, like systematics arising from LPF self-gravity, or the Newtonian background, are examined and shown not to be a problem. We also investigate the impact of a negative observational result upon the free-function determining the theory. We demonstrate that, if Newton's gravitational constant is constrained not be re-normalized by more than a few percent, only contrived MONDian free-functions would survive a negative result. There are exceptions, e.g. free-functions not asymptoting to 1 in the Newtonian limit, but rather diverging or asymptoting to zero (depending on their mother relativistic MONDian theory). Finally, we scan the structure of all proposed relativistic MONDian theories, and classify them with regards to their non-relativistic limit, finding three broad cases (with a few sub-cases depending on the form of the free function). It is appears that only the Einstein-Aether formulation, and the sub-cases where the free-function does not asymptote to 1 in other theories, would survive a negative result without resorting to ``designer'' free-functions. ",
        "introduction": "Einstein's theory of General Relativity (GR) and the $\\Lambda$CDM standard model are two cornerstones of modern cosmology. They posit that the gravitational effects of large scale structures in the universe (such as galaxies and clusters of galaxies) cannot be explained by luminous, baryonic matter alone, but rather that an additional cold (pressureless) and dark (non-luminous) matter component is needed. However, in the absence of direct observational evidence for dark matter, it remains nothing but a useful calculational device. For as long as this is true, it is scientifically healthy to explore alternative explanations for the anomalous gravitational dynamics, namely by modifying the theory of gravity itself.  MOdified Newtonian Dynamics (or MOND~\\cite{Milgrom:1983ca}) is one such scheme valid in the non-relativistic regime. It was first proposed to explain observed dynamical properties of galaxies without invoking dark matter. More recently it has been incorporated into relativistic theories~\\cite{BSTV,aether,aether1,Milgrom:2009gv,Milgrom:2010cd}, following from the ground breaking proposal of ``TeVeS'' by Bekenstein~\\cite{teves}. Relativistic extensions are needed for reasons beyond logical completeness: they are required to explain, for example, phenomenology associated with lensing and cosmology~\\cite{kostasrev},  where dark matter is also usually employed. When the whole picture is assembled the conflict between MOND and dark matter leaves considerable scope for doubt over the interpretation of new astrophysical and cosmological data. A fair comparison requires re-evaluating, within each approach, the whole set of assumptions underlying the new observations. For this reason the debate would benefit from a direct probe, in the form of a laboratory or Solar System experiment. This has been proposed in various forms, namely in planetary data~\\cite{Blanchet} appealing to the exterior field effect~\\cite{Milgromss}.  The fact that MOND predicts anomalously strong tidal stresses in the vicinity of saddle points of the Newtonian potential has been advocated as one such  decisive direct test~\\cite{Bekenstein:2006fi}. The forthcoming LISA Pathfinder mission~\\cite{LISA} presents the perfect opportunity for its realization, as a preliminary feasibility study has demonstrated~\\cite{Bevis:2009et,companion}. The purpose of this paper is to provide a detailed quantitative evaluation of the power of a MONDian saddle test using LPF, predicated upon a scenario where a mission extension is granted. The extension would involve redirecting the spacecraft from Lagrange point L1 to a saddle of the Earth-Moon-Sun system~\\cite{companion} once its nominal mission at L1 is completed. In establishing the scientific case our efforts in this paper are twofold.  In the first part of the paper we propose some basic data analysis tools and evaluate their expected performance. These tools are an adaptation of the ``noise-matched filters'' employed in gravitational wave detection~\\cite{sathya}. Their implementation benefits from a major simplification: for a saddle test we do know the template's starting point in time. A number of pitfalls and potential systematics found in detection of gravitational waves are therefore expected to be absent. The filter's optimal signal to noise ratio (SNR) allows us to quantify with a single number the predicted outcome for any experiment. Assuming a ``standard'' MONDian theory, the unknowns reduce to the instrument performance (the noise properties) and the trajectory past the saddle (its impact parameter). For each of these we can condense the expected outcome of a LPF test in a single number: the forecast SNR assuming MOND is correct.   Our central results are in Section~\\ref{secsnr}, particularly in Figs.~\\ref{fig:SNR contours} and~\\ref{fig:SNR-improved}, where the optimal SNR is plotted against noise level and saddle impact parameter. In a cataclysmic scenario for instrument performance and saddle approach we'd still achieved ${\\rm SNR}\\approx 5$. For less pessimistic assumptions, high double figures are easily reached. We examine the effect of the spacecraft speed as it flies past the saddle, showing that just about any typical speed will turn out to be optimal. This is due to a remarkable coincidence, spelled out in Section~\\ref{systematics} and in the concluding section of this paper. In Section~\\ref{systematics} we also show that possible systematic errors, such as self-gravity or the Newtonian background, are in fact harmless.  In the second part of this paper, and complementing the work just described, we spell out the generality, or otherwise,  of the conclusions in the first part, and examine the implications of a negative observational result. Just how comprehensively would the failure to detect the predicted high SNRs rule out the MONDian paradigm as a whole? As explained in Section~\\ref{theory} the large menagerie of proposed relativistic MONDian theories practically all reduce to the same non-relativistic limit as TeVeS, and virtually all theories fall into 3 categories. One then is left with a free-function, $\\mu$, and the question is, how much leeway does it provide for evading a negative result? In Section~\\ref{frefn} we review previously proposed free functions, rewriting them under a unified notation. We then lay down conditions for what should be permissible {\\it simple} free functions at the most basic level (by simple we mean with a minimum number of regimes and scales). Briefly we require that: (I) The theory shouldn't renormalize Newton's gravitational constant by more than a few percent in the Newtonian regime; (II) The theory should predict the usual MONDian effects when the Newtonian acceleration drops below acceleration scale $a_0$. These constraints were implemented in TeVeS and set the standard for a viable theory with useful astronomical implications. We show that {\\it all natural functions satisfying these conditions result in similar SNRs for an LPF saddle test, as long as the impact parameter is smaller than $400\\unit{km}$}. The only exceptions are free-functions with a divergence (and with the rest of their domain excised) such as those suggested in~\\cite{Angus,FamaeyN}; however these {\\it may} fall foul of existing Solar system constraints and other requirements~\\cite{Fam-gaugh,FamaeyN,ssconst,Gentile}. We leave for a future publication a more detailed examination of these functions.    Of course one may open the doors to free-functions with more structure: e.g. functions with three or more regimes instead of the minimal two. In Section~\\ref{nullres} we quantify how contrived the free function $\\mu$ would have to be, for the theory to survive a negative result. We find that, for bounded functions, only a $\\mu$ turning from 1 (Newtonian regime) into an intermediate power-law, $\\mu\\propto z^n$, and only then into the MONDian $\\mu\\propto z$, would be viable. The intermediate $n$ would have to be very different from 1 even with undemanding requirements on impact parameter and noise. Another possibility are the free-functions suggested in~\\cite{Zhao,Bruneton,McGaughMW,Zhaoeco}, for which galaxy rotation fits and Solar system constraints are combined to motivate more structured free-functions. All of these invoke three regimes (and so have two implied scales). Even then, in Section~\\ref{nullres1} we show that these would be within striking distance for a saddle test with LPF. Thus, with the exception of diverging $\\mu$, only ``designer'' $\\mu$ would bypass a negative result. Although this conclusion is derived for TeVeS-like theories it might be more general. With the usual honorable exceptions and provisos, the Einstein-Aether formulation~\\cite{aether,aether1} and a particular case of Milgrom's bimetric theory~\\cite{Milgrom:2009gv,Milgrom:2010cd} might be the only relativistic realizations of MOND to realistically survive a negative result from a saddle test, a matter we'll examine in a future publication.  We conclude that a LPF test has both the power to detect MOND with a high SNR should it be true, and to  rule it out, should a negative result be obtained. ",
        "conclusions": "\\label{concs} In this paper we showed how a LPF saddle flyby would either detect MOND to a high SNR or rule it out, if not comprehensively, at least to a large extent. The former conclusion could be expected. Even though in Sections~\\ref{secsnr} and~\\ref{systematics} we provided rigorous and quantitative SNR estimates, the high levels forecast can be understood with a ``back of the envelope'' calculation. The exercise highlights an uncanny coincidence. The accelerometer aboard LPF has a non-white noise profile, dipping in the region of the mHz, i.e. in the rough time scale of minutes. The motivation for this design lies in the gravitational wave signals to be targeted by LISA. It just so happens that the MONDian bubbles of anomalous tidal stresses around the Earth-Sun-Moon saddles have a length scale of the order of a hundred kilometers. Anything free-falling in the Earth-Moon region has a typical speed of the order of 1km/s. Thus, the time scale for crossing a MONDian bubble will be of the order of minutes: just where the instrument performance is optimal. This is a remarkable coincidence. Scribbling on the back of an envelope, using the expression for the SNR of a noise-matched filter and the order of magnitude of the stresses and noise, promptly reveals double figure SNRs.  The question then arises as to how generic this conclusion is, or conversely, should a negative result be found, how thoroughly have we ruled out MOND. In Section~\\ref{theory} we scanned the full array of MONDian theories, showing that in the non-relativistic regime they fall into only 3 categories  (which we labelled type I, IIA/B and III). Even if we restrict ourselves to type I theories (a class containing the vast majority of these theories, including TeVeS), we benefit from the leverage of a free function, $\\mu$. Could MONDologists use $\\mu$ to survive a negative result?  In Sections~\\ref{frefn} we examined $\\mu$ functions on offer in the literature and laid down criteria for reasonable $\\mu$ based on astrophysical usefulness, viability in the face of constraints, {\\it and naturalness}. We found that once these criteria are taken into account the size of the MOND bubble, $r_0$, is fixed. There are exceptions to this rule; e.g. diverging $\\mu$ functions, but these {\\it may} be subject to other constraints (they will be examined elsewhere). Predictions for what happens inside the bubble are also model independent; however the tidal stress anomalies outside the bubble depend on the transient from MONDian into Newtonian regime, with a fall-off which is indeed model dependent. Thus, for impact parameters smaller than $r_0$ the predicted SNRs are robust, and do not change substantially with the model. For the currently expected $b$ (around $50\\unit{km}$, with $r_0\\sim 380\\unit{km}$) this is indeed the case.  Therefore the only way for MOND to wriggle out of a negative LPF result would be to change the bubble size $r_0$. This can only be accomplished with ``designer'' $\\mu$-functions. If $\\mu$ is allowed to have two scales and two power-laws away from its Newtonian value of 1, then it is possible to bypass a negative LPF result. Even for undemanding noise levels and impact parameters, the intermediate power becomes very contrived. Similarly fine-tuned functions have been proposed in the literature. In Section~\\ref{nullres1}, we showed how LPF could be used to constrain them. The point remains that one would have to bend backwards to accommodate a negative result, within type I theories. There are exceptions to this rule (for example diverging $\\mu$ functions) and these derserve further attention.   Although we didn't present quantitative results for type II theories, the same conclusions apply qualitatively to type IIB theories (but not type IIA theories). If $G$ is renormalized and the free function $\\nu$ is chosen to produce the same phenomenology as type I theories (in particular with regards to $G_{Ren}$ and MONDian behaviour), then the MOND bubble has the same size, and the anomalous tidal stresses are of the same order.  As explained in Section~\\ref{theory}, in both types of theory MONDian behaviour is due to an extra field $\\phi$, and if one attends simultaneously to $G_{Ren}\\approx G$ and $\\phi\\sim \\Phi_N$ for $a_N\\approx a_0$, then MONDian behavior in $\\phi$ should be triggered at the same Newtonian acceleration $a_N=a_N^{trig}\\gg a_0$. This implies a MONDian bubble of the same size $r_0$. Furthermore the (also $\\nu$-independent) effects inside the bubble are different from type I predictions, but stronger. Type II theories don't have a curl field (in the sense define above), a feature which softens the anomalous tidal stresses in type I theories~\\cite{Bekenstein:2006fi}. A detailed quantitative prediction for type II theories is currently being investigated (see~\\cite{aliqmond}; also we reference here a paper on the matter which has appeared since the present paper was submitted~\\cite{typeII}.)   However, it may be that the relativistic mother theory is set up in such a way that the cosmological and non-relativistic $G$ coincide, the case of type IIA theories. In this case the MOND bubble around the saddle is very small. Likewise type III theories (or, rather, its single relativistic realization, the Einstein-Aether theory) produce effects around saddles which are unobservable with current technology. In these theories $G$ is not renormalized and $a^{trig}=a_0$, so that the MOND bubble is a few meters across. Remarkably, solar system tests are extremely constraining upon type III theories, due to the so-called external field effect~\\cite{Blanchet}.  By contrast solar system effects for type I and IIB theories are suppressed by a factor of $\\kappa/4\\pi$. Thus saddle tests and planetary orbits seem to be complementary in constraining MONDian theories.  We close by noting that we could, of course, detach our considerations entirely from the MOND paradigm (as an alternative to dark matter), and regard these theories formally as a class on alternative theories of gravity (see~\\cite{Clifton11} for an extensive review). It is remarkable that only three classes of theories emerge in the non-relativistic regime, which we labelled type I, II and III in Section~\\ref{theory}. We could then view $\\kappa$ and $a_0$ as free parameters, converting a LPF saddle flyby into a constraint or a detection in this space. We are currently working on this alternative approach."
    },
    "1107/1107.1243_arXiv.txt": {
        "abstract": "We study the change in cosmic-ray pressure, the change in cosmic-ray density, and the level of cosmic-ray induced heating via \\alfven-wave damping when cosmic rays move from a hot ionized plasma to a cool cloud embedded in that plasma.  The general analysis method outlined here can apply to diffuse clouds in either the ionized interstellar medium or in galactic winds.  We introduce a general-purpose model of cosmic-ray diffusion building upon the hydrodynamic approximation for cosmic rays (from McKenzie \\& V\\\"olk and Breitschwerdt and collaborators).  Our improved method self-consistently derives the cosmic-ray flux and diffusivity under the assumption that the streaming instability is the dominant mechanism for setting the cosmic-ray flux and diffusion.  We find that, as expected, cosmic rays do not couple to gas within cool clouds (cosmic rays exert no forces inside of cool clouds), that the cosmic-ray density does not increase within clouds (it may slightly decrease in general, and decrease by an order of magnitude in some cases), and that cosmic-ray heating (via \\alfven-wave damping and not collisional effects as for $\\sim 10$\\,\\MeV cosmic rays) is only important under the conditions of relatively strong (10\\,\\muG) magnetic fields or high cosmic-ray pressure ($\\sim 10^{-11}$\\,ergs\\,cm$^{-3}$). ",
        "introduction": "The possibility of a rise or fall in cosmic-ray density within diffuse clouds and/or molecular cores has been studied over many years by researchers interested either in gamma-ray emission from cool clouds \\citep[e.g.,][]{Skilling1971, SkillingStrong1976, CesarskyVoelk1977, StrongSkilling1977, CesarskyVoelk1978, Morfill1982a, Morfill1982b, Voelk1983} or in ionization by cosmic rays within clouds \\citep[e.g.,][]{HartquistEtAl1978, SuchkovEtAl1993, PadoanScalo2005, PadovaniEtAl2009, Papadopoulos2010}.  Another regime where this interaction may be important is in cool clouds that are readily observed within galactic winds \\citep[e.g.,][]{HeckmanEtAl1990, Martin2005, RupkeEtAl2005, SteidelEtAl2010}.  In the context of such winds, cosmic rays are known to help heat and drive outflows of fully ionized gas \\citep{BreitschwerdtEtAl1991, EverettEtAl2008, SocratesEtAl2008, EverettEtAl2010}; can they help heat and accelerate cool clouds embedded in a hot wind?  We will examine this problem using a set of hydrodynamic equations that treat both convective and diffusive transport of cosmic rays self-consistently; this has rarely been done in studies of cosmic-ray dynamics, and has never been done in the context of cool clouds. Before we start, however, it is important to understand the variety of previous work on this problem.  We start, therefore, with a review of models of cosmic-ray interaction with cool clouds in Section~\\ref{reviewSection}; we cover models of cosmic-rays in the ISM (Section~\\ref{ismStudiesReview}) \\& cosmic rays in galactic winds (Section~\\ref{galacticWindsReview}), followed by an introduction to our method (Section~\\ref{ourModel}).  Section~\\ref{cosmicRayHydro} defines our equations for cosmic-ray hydrodynamics, and Section~\\ref{cloudModels} applies those models to a simple hot-gas/cool-cloud interface, and examines how robust those results are.  Finally, Section~\\ref{conclusions} summarizes our results and compares \\& contrasts those results with previous work. ",
        "conclusions": ""
    },
    "1107/1107.1769_arXiv.txt": {
        "abstract": "We employ a three-dimensional (3D) reconstruction technique, for the first time to study the kinematics of six coronal mass ejections (CMEs), using images obtained from the COR1 and COR2 coronagraphs on board the twin STEREO spacecraft, as also the eruptive prominences (EPs) associated with three of them using images from the Extreme UltraViolet Imager (EUVI). A feature in the EPs and leading edges (LEs) of all the CMEs was identified and tracked in images from the two spacecraft, and a stereoscopic reconstruction technique was used to determine the 3D coordinates of these features. True velocity and acceleration were determined from the temporal evolution of the true height of the CME features. Our study of kinematics of the CMEs in 3D reveals that the CME leading edge undergoes maximum acceleration typically below $2\\Rsun$. The acceleration profiles of CMEs associated with flares and prominences exhibit different behaviour. While the CMEs not associated with prominences show a bimodal acceleration profile, those associated with prominences do not. Two of the three associated prominences in the study show a high and rising value of acceleration up to a distance of almost $4\\Rsun$ but acceleration of the corresponding CME LE does not show the same behaviour, suggesting that the two may not be always driven by the same mechanism. One of the CMEs, although associated with a C-class flare showed unusually high acceleration of over $1500\\mpss$. Our results therefore suggest that only the flare-associated CMEs undergo residual acceleration, which indicates that the flux injection theoretical model holds good for the flare-associated CMEs, but a different mechanism should be considered for EP-associated CMEs. ",
        "introduction": "\\label{S:intro}  Coronal mass ejections (CMEs) result from a loss of equilibrium in the magnetic configuration in the solar corona \\citep{Priest1988,Klimchuk2001}. Several factors like, flux emergence, flux cancellation, reconnection, shear, etc., are thought to be responsible for this loss of equilibrium \\citep{Forbes.etal2006,Seaton.etal2011}. Once the equilibrium is lost, the energy needed by the CME for its propagation is derived from the surrounding magnetic field \\citep{Forbes2000,Low2001}. Very often, the energy of the surrounding field is sufficient not only to propel a CME, but also to accelerate it \\citep{Alexander2006}.  \\citet{Zhang.Dere2006} have categorised the evolution of CMEs into a three-phase process involving initiation, acceleration and propagation (Figure 1 in their paper). According to \\citet{Zhang.Dere2006}, the initiation phase is the phase of slow rise of CMEs, and in the acceleration phase they undergo a very rapid increase in their velocity, while, in the propagation phase, the CME velocity remains more or less constant, \\thatis, it experiences almost zero acceleration. Using LASCO \\citep{Brueckner.etal1995} observations on board the SoHO spacecraft \\citep{Domingo.etal1995a}, \\citet{Yashiro.etal2004} have observed that the CME velocity in the outer corona varies from less than $100\\kmps$ to over $3000\\kmps$. The propagation of CMEs can be understood if we consider the forces acting on them, which are the Lorentz force, gravitational force, and drag because of the ambient solar wind. Of the three forces, the drag force is the strongest beyond a few solar radii, and the other two can be neglected \\citep{Gopalswamy.etal2001b,Cargill2004,Vrsnak.etal2010}. This is further supported by results obtained by \\citet{Gopalswamy.etal2000}. They have observed that although initial CME speeds range from $124-1056\\kmps$, the speeds of the corresponding interplanetary ejecta are found to lie in the range of $320-650\\kmps$, which is more or less the speed of the ambient solar wind. \\citet{Cargill2004} have reported that speeds of interplanetary CMEs (ICMEs) corresponding to CMEs with speeds ranging from $100-2000\\kmps$, as measured from coronagraphs, lie within $100-200\\kmps$ of the ambient solar wind. However, the time a CME takes to reach the Earth, the transit time, is known to vary from less than a day to over four days. This indicates that most of the CME dynamics occurs closer to the Sun. \\citet{Vrsnak.etal2010} have reported that transit times of broad, low-mass CMEs depend mainly on the surrounding solar wind speed, while those of narrow, massive CMEs depend mainly on the initial speeds of the CMEs. Recently, \\citet{Manoharan.Rahman2011} have also found that most of the ICMEs tend to attain speeds close to that of the ambient solar wind, and have estimated travel times of the CMEs to reach a distance of \\mbox{1\\,AU} based on the CME initial speed and drag due to solar wind.  CMEs have been classified on the basis of their source regions. \\citet{Gosling.etal1976}, using the coronagraph on Skylab spacecraft, were the first ones to report that CMEs associated with flares are faster than those associated with prominences. This was supported by observation of CMEs by \\citet{MacQueen.Fisher1983} who used the \\textit{K}-coronameter at Mauna Loa Solar Observatory. In addition, they also observed that the former type showed smaller acceleration with increase in height than the latter. \\citet{Sheeley.etal1999} have also reported a similar result based on their technique to track features observed in SoHO/LASCO coronagraphs. \\citet{Moon.etal2002} in a statistical study involving over 3200 CMEs observed from SoHO/LASCO have reported that flare-associated CMEs have a higher median speed than those associated with EPs. Their study also found that although the median acceleration of all the events is zero, it decreases a little for CMEs with high speeds ($>500\\kmps$). \\citeauthor{Srivastava.etal1999a} (\\citeyear{Srivastava.etal1999a}; \\citeyear{Srivastava.etal2000}) have found that gradual CMEs attain the speed of the ambient solar wind at about $20\\Rsun$ from the Sun. Results from \\citet{Gopalswamy.etal2001b} also are consistent with this study, who reported deceleration as high as $-100\\mpss$ for fast CMEs (speed $>900\\kmps$) from a combined study of SoHO/LASCO and radio observations from Wind spacecraft.  \\sloppy \\citet{Chen.Krall2003} have studied acceleration of three CMEs using SoHO/LASCO observations, and proposed that CME acceleration occurs in two phases, the `main' phase and the `residual' phase. While most of the acceleration occurs in the main phase, there lies a second phase of acceleration known as the residual acceleration in the outer corona. \\textbf{\\citet{Chen.Krall2003} and \\citet{Chen.etal2006} have identified the main acceleration phase as the interval over which Lorentz force is the most dominant, while during the residual phase, Lorentz force is comparable to the two other force, viz, gravity and drag.} They have employed the magnetic flux rope model \\citep{Chen1989} to show a relation between the height at the peak of main acceleration phase, and the  footpoint separation of the CME flux rope. In their model, \\citet{Chen.Krall2003} have proposed that a change in duration of the flux injection \\citep{Krall.etal2000} determines the strength of the residual acceleration phase. Similarly, \\citet{Zhang.Dere2006} have also reported two such phases of acceleration based on their study of 50 CMEs observed from SoHO/LASCO.  All the studies cited above use a single viewpoint to observe the CMEs. The results then inherently suffer from projection effects of the transients on to the plane of the sky. In order to overcome this, we decided to look at CMEs from the stereoscopic vision of Solar TErrestrial RElations Observatory (STEREO) \\citep{Kaiser.etal2008}. The STEREO spacecraft provide two viewpoints of the prominences and the associated CMEs. We have used a stereoscopic reconstruction technique to determine the true physical coordinates of a solar feature \\citep{Joshi.Srivastava2011}. The stereoscopic reconstruction would allow us to observe evolution of the true height of prominences and CMEs, and hence their true velocity and acceleration. From this we can examine if the acceleration truly exhibits bimodal profile as the model suggests. This will also give us a clue about the initiation and propagation of CMEs in the corona. ",
        "conclusions": "\\label{S:conclude}  We have analysed six CMEs from the coronagraphs COR1 and COR2, and the associated EPs in three of the cases from EUVI on board the identical STEREO\\,A and B spacecraft. We identified and tracked a feature in the LE of all the CMEs in both COR1 and COR2, and in the associated prominences, wherever applicable. While most of the earlier studies on CME acceleration were carried out using projected measurements, we have used a stereoscopic reconstruction technique \\citep{Joshi.Srivastava2011} to obtain the true coordinates, and hence the true speed and acceleration of the feature. On fitting a polynomial function to the true height, the speed and acceleration of the CMEs as a function of time and true height were determined. The results of the kinematic study of EPs and the CME LEs are shown in Figures~\\ref{F:res16nov}--\\ref{F:res01aug}. We summarise the results obtained from the reconstruction in Table~\\ref{T:summa}.   \\vspace{-0.2cm} \\begin{center} \\begin{table}[!tbp] \\begin{tabular}{lrrrrrr} \\hline Event & $v_{max}$ & height of & $a_{max}$ & height of & $v$ at & $a$ at \\\\ & (\\kmps) & $v_{max}$ ($\\Rsun$) & (\\mpss) & $a_{max}$ ($\\Rsun$) & 10$\\Rsun$ & 10$\\Rsun$ \\\\ \\hline 16 Nov 2007 LE & 451  & 12.2  & 50    & 2.2  & 408  & 16   \\\\ 31 Dec 2007 LE & 876  & 13.0  & 1524  & 1.9  & 860  & 2    \\\\ ~9 Apr 2008 LE & 533  & 7.6   & 123   & 2.3  & 488  & -15  \\\\ 16 Dec 2009 LE & 356  & 5.8   & 90    & 1.9  & 488  & -15  \\\\ 13 Apr 2010 LE & 522  & 12.6  & 61    & 1.9  & 193  & 36   \\\\ ~1 Aug 2010 LE & 567  & 4.4   & 213   & 4.4  & ---  & ---  \\\\ ~9 Apr 2008 EP & 268  & 3.5   & 104   & 1.2  & ---  & ---  \\\\ 13 Apr 2010 EP & 377  & 3.9   & 141   & 3.9  & ---  & ---  \\\\ ~1 Aug 2010 EP & 224  & 2.9   & 34    & 1.6  & ---  & ---  \\\\ \\hline \\end{tabular} \\caption{Summary of the 6 LEs and 3 EPs analysed using three-dimensional reconstruction. $v_{max}$ and $a_{max}$ denote the maximum speed and acceleration of the CME calculated. The heights at which CMEs attained these values are also provided. The last two columns show the speed and acceleration of the CMEs at a distance of $10\\Rsun$.}\\label{T:summa} \\end{table} \\end{center} \\vspace{-1.0cm}  It is believed that most of the CME acceleration typically occurs in the lower corona. \\citet{Chen.Krall2003} have found the height of maximum acceleration of CME to be $2-3\\Rsun$ from a study of several CMEs,  \\citet{Vrsnak2001a} have considered this height to be $4\\Rsun$. However, from our reconstructed results (Figures~\\ref{F:res16nov}--\\ref{F:res01aug}), we observe that in all the cases studied here, the peak of main phase of acceleration lies below the true height of $2\\Rsun$. This indicates that most of the CME dynamics occurs closer to the Sun than previously believed, as shown by \\citet{Chen.etal2006} from a comparison of observations and models.  Earlier studies \\citep{Zhang.etal2001,Chen.Krall2003}, have observed CMEs in all the three SoHO/LASCO coronagraphs which together cover a range from $1.1-32\\Rsun$. In such studies, the initiation phase of the CME, as well as peak of their acceleration could be observed. The COR1 and COR2 coronagraphs together image the solar corona from $1.4$ to $15.0\\Rsun$, however these are only the plane-of-sky FOVs of the coronagraphs. The minimum value of true height of reconstructed features the corona is approximately $2.0\\Rsun$. Thus, in most of the cases we do not capture the rise phase acceleration of the LE of CME. In all but one case studied here, the acceleration peak has already passed from the time we start observing the CME. At this point, it is necessary to point out that the heights determined in this study are true heliocentric distances, hence they are seen to be significantly different than the heights obtained from previous studies which relied upon observations from a single spacecraft.  The CME on 2007 December 31 was associated with a flare having X-ray class C8, however it still showed a very high value of acceleration of over $1500\\mpss$. Earlier studies have shown that the acceleration phase of CMEs coincides with the increase in soft X-ray flux due to the associated flare \\citep{Neupert.etal2001,Shanmugaraju.etal2003}. \\citet{Maricic.etal2007} have also shown that both the velocity and acceleration of the CME show a significant correlation with the X-ray class of the associated flare. As per the least squares fit obtained from their study, acceleration of the CME associated with a C8 flare should be around $300\\mpss$. The value calculated by us however, is 5 times more, suggesting that the flare energy alone might not be the only one to drive the CMEs. In such a scenario, the supposition that impulsive and gradual CMEs are respectively associated with flares and EPs \\citep{MacQueen.Fisher1983,Moon.etal2002} should also be subjected to further scrutiny. Also, deviations to the findings reported by \\citet{Maricic.etal2007}, where acceleration of CMEs is correlated with the X-ray class of the associated flare, should not be ignored. \\citet{Raftery.etal2010} have used soft and hard X-ray observations in addition to STEREO observations \\citep{Lin.etal2010} to analyse the 2007 December 31 CME, and have found that it follows the tether-cutting reconnection model.  The results obtained from reconstruction were used to determine maximum acceleration, average acceleration, acceleration magnitude, and acceleration duration, attained by the CMEs and EPs. The time interval between the maximum and the zero value of acceleration is termed as the acceleration duration, while the velocity increase during this time divided by acceleration duration is termed as the acceleration magnitude \\citep{Zhang.Dere2006}. The very high value of acceleration for the 2007 December 31 CME makes the event an `outlier', hence, we have not included that value in the scatter plots in Figure~\\ref{F:max_spd}. The left panel shows speed at maximum acceleration plotted against the maximum acceleration. From this figure, we find that in the events studied by us, higher the maximum value of acceleration, higher is the speed at that instant. While, the right panel of Figure~\\ref{F:max_spd} shows scatter plot of maximum acceleration and height at the instance of maximum acceleration. This scatter plot suggests, that higher the acceleration, higher up in the corona it occurs.  \\textbf{Although the acceleration in this study is determined up to the COR2 FOV, it may not be the value with which the CME is travelling at larger distances from the Sun. Based on interplanetary measurements, it was shown earlier by \\citet{Gopalswamy.etal2000} that slow CMEs tend to accelerate, while the faster ones tend to decelerate. Recently, \\citet{Davis.etal2010} have measured the speeds of 26 CME from the Heliospheric Imagers (HI), which are part of the SECCHI suite on the STEREO spacecraft. In their study, they have found that CMEs with speeds less than $400\\kmps$ in COR2 FOV have higher speeds in HI FOV, and vice-versa. Thus they have cautioned that a CME may undergo genuine acceleration even in the HI FOV, which extend from $15-84\\Rsun$ for HI-1 and $66-318\\Rsun$ for HI-2.}  Previous studies have reported that acceleration of a CME shows bimodal distribution \\citep{Chen.Krall2003}. We observe such a bimodal distribution in 3 CMEs, the ones which are not associated with prominence eruptions. The residual acceleration for the very impulsive 2007 December 31 CME was $90\\mpss$, while for the CMEs on 2007 November 16 and 2009 December 16, it was found to be $18\\mpss$ and $-2\\mpss$, respectively. The other CMEs, which are associated with prominences do not show such an acceleration profile. \\citet{Chen.Krall2003} have invoked the flux injection mechanism to trigger an eruption in a magnetic flux rope, which leads to the residual acceleration phase. In the cases analysed here, we observe that only the flare-associated CMEs undergo residual acceleration, which indicates that flux injection seems to be a good explanation for eruption of the flare-associated CMEs studied here, but a different mechanism should be considered for EP-associated CMEs.  Of the three CMEs associated with prominences, the 2010 April 13 and 2010 August 1 were associated with large quiescent polar crown prominences, while the one on 2008 April 9 was associated with an active-region prominence. We find that the prominences on 2008 April 9 and 2010 April 13 showed a strong positive acceleration in the COR1 FOV, when their heights were close to almost $4\\Rsun$. During the same time however, acceleration of the CME LE was decreasing. This indicates that even at a height of $4\\Rsun$, forces acting on the CME and the EP cannot be considered to be the same, as suggested by \\citet{Srivastava.etal2000} and \\citet{Maricic.etal2004}.  Thus, in this study, from the 3D reconstruction of six CMEs and EPs associated with three of them, we have observed some aspects of their acceleration, as detailed above, which were not previously reported. We find that the maximum CME acceleration occurs at a height of less than $2\\Rsun$, where earlier, this height was believed to be between $2-4\\Rsun$. The bimodal acceleration profile was not observed in EP-associated CMEs, but in only those CMEs that were not associated with EPs. Two of the three prominences in the study showed a high and rising value of acceleration at a distance of almost $4\\Rsun$ but the corresponding CME LE does not show the same behaviour. The CME on 2007 December 31, showed acceleration of over $1500\\mpss$, which is unusually high for a CME associated with a C-class flare.   The authors thank the STEREO/SECCHI consortium for providing the data. The SECCHI data used here were produced by an international consortium of the Naval Research Laboratory (USA), Lockheed Martin Solar and Astrophysics Lab (USA), NASA Goddard Space Flight Center (USA), Rutherford Appleton Laboratory (UK), University of Birmingham (UK), Max-Planck-Institut for Solar System Research (Germany), Centre Spatiale de Li$\\grave{\\textrm e}$ge (Belgium), Institut d'Optique Theorique et Appliqu$\\acute{\\textrm e}$e (France), Institut d'Astrophysique Spatiale (France). Work by N.S. partially contributes to the research on collaborative NSF grant \\mbox{ATM-0837915} to Helio Research.  \\begin{figure} \\centering \\includegraphics[width=0.30\\textwidth,clip=]{c1_b_16nov07_0945.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c1_a_16nov07_0945.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c2_b_16nov07_1507.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c2_a_16nov07_1507.eps} \\\\ \\caption{Images of the CME on 2007 November 16 seen in images from COR1 (top panels) and COR2 (bottom panels), as seen from STEREO\\,B (left) and A (right).}\\label{F:img16nov} \\end{figure}   \\begin{figure}[!p] \\centering \\includegraphics[width=0.30\\textwidth,clip=]{c1_b_31dec07_0115.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c1_a_31dec07_0115.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c2_b_31dec07_0207.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c2_a_31dec07_0207.eps} \\\\ \\caption{Images of the CME on 2007 December 31 seen, similar to Figure~\\ref{F:img16nov}.}\\label{F:img31dec} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=0.30\\textwidth,clip=]{eu_b_09apr08_1006.eps} \\includegraphics[width=0.30\\textwidth,clip=]{eu_a_09apr08_1006.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c1_b_09apr08_1055.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c1_a_09apr08_1055.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c2_b_09apr08_1308.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c2_a_09apr08_1308.eps} \\\\ \\caption{Images of the EP and associated CME on 2008 April 9, seen in images from EUVI 304\\Ang\\ (top panels), COR1 (middle panels) and COR2 (bottom panels), as seen from STEREO\\,B (left) and A (right).}\\label{F:img09apr} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=0.30\\textwidth,clip=]{c1_b_16dec09_0225.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c1_a_16dec09_0225.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c2_a_16dec09_0454.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c2_b_16dec09_0454.eps} \\\\ \\caption{Images of the CME on 2009 December 16 seen, similar to Figure~\\ref{F:img16nov}.}\\label{F:img16dec} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.30\\textwidth,clip=]{eu_b_13apr10_0806.eps} \\includegraphics[width=0.30\\textwidth,clip=]{eu_a_13apr10_0806.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c1_b_13apr10_0945.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c1_a_13apr10_0945.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c2_b_13apr10_1254.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c2_a_13apr10_1254.eps} \\\\ \\caption{Images of the EP and the associated CME on 2010 April 13, similar to Figure~\\ref{F:img09apr}.}\\label{F:img13apr} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=0.30\\textwidth,clip=]{eu_b_01aug10_0746.eps} \\includegraphics[width=0.30\\textwidth,clip=]{eu_a_01aug10_0746.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c1_b_01aug10_0925.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c1_a_01aug10_0925.eps} \\\\ \\includegraphics[width=0.30\\textwidth,clip=]{c2_b_01aug10_1024.eps} \\includegraphics[width=0.30\\textwidth,clip=]{c2_a_01aug10_1024.eps} \\\\ \\caption{Images of the EP and the associated CME on 2010 August 1, similar to Figure~\\ref{F:img09apr}.}\\label{F:img01aug} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.465\\textwidth,clip=]{rcn_2007nov16_spac1.eps} \\includegraphics[width=0.470\\textwidth,clip=]{rcn_2007nov16_vsht1.eps} } \\caption{Results from the stereoscopic reconstruction applied to a feature in LE of the CME on 2007 November 16. Left: True height of the CME feature against time in the top panel, followed by the true speed and acceleration against time in the middle and bottom panels. Right: True height of the CME against time at the top, followed by the true speed and acceleration against true height in the middle and bottom panels. Plus signs (+) and asterisks (*) indicate that the feature was observed in COR1 and COR2 FOVs, respectively.}\\label{F:res16nov} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.46\\textwidth,clip=]{rcn_2007dec31_spac1.eps} \\includegraphics[width=0.46\\textwidth,clip=]{rcn_2007dec31_vsht1.eps} } \\caption{Results from stereoscopic reconstruction for CME on 2007 December 31, similar to Figure~\\ref{F:res16nov}}\\label{F:res31dec} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.465\\textwidth,clip=]{com_2008apr09_pL_spac.eps} \\includegraphics[width=0.470\\textwidth,clip=]{com_2008apr09_pL_vsht.eps} } \\caption{Results from stereoscopic reconstruction for CME on 2008 April 9, similar to Figure~\\ref{F:res16nov}. Here, triangles ($\\vartriangle$) indicate feature observed in EUVI FOV.}\\label{F:res09apr} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.465\\textwidth,clip=]{rcn_2009dec16_spac1.eps} \\includegraphics[width=0.470\\textwidth,clip=]{rcn_2009dec16_vsht1.eps} } \\caption{Results from stereoscopic reconstruction for CME on 2009 December 16, similar to Figure~\\ref{F:res16nov}}\\label{F:res16dec} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.465\\textwidth,clip=]{com_2010apr13_pL_spac.eps} \\includegraphics[width=0.470\\textwidth,clip=]{com_2010apr13_pL_vsht.eps} } \\caption{Results from stereoscopic reconstruction for CME on 2010 April 13, similar to Figure~\\ref{F:res09apr}}\\label{F:res13apr} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.465\\textwidth,clip=]{com_2010aug01_pL_spac.eps} \\includegraphics[width=0.470\\textwidth,clip=]{com_2010aug01_pL_vsht.eps} } \\caption{Results from stereoscopic reconstruction for CME on 2010 August 1, similar to Figure~\\ref{F:res09apr}}\\label{F:res01aug} \\end{figure}   \\begin{figure} \\centerline { \\includegraphics[width=0.48\\textwidth,clip=]{maxacc_spd_a.eps} \\includegraphics[width=0.47\\textwidth,clip=]{maxacc_hgt_a.eps} } \\caption{Two scatter plot showing speed (left panel), and height (right panel) at the instance of maximum acceleration and the maximum acceleration itself of the 6 CMEs and 3 EPs studied. The legend shows data points corresponding to each event. Data point for 2007 December 31 is not shown since it has a very high value of maximum acceleration.}\\label{F:max_spd} \\end{figure}"
    },
    "1107/1107.5656_arXiv.txt": {
        "abstract": "{ It has recently been shown that the terrestrial planets and asteroid belt can be reproduced if the giant planets underwent an inward-then-outward migration (the \"Grand Tack\"; Walsh et al 2011).  Inward migration occurs when Jupiter opens a gap and type II migrates inward.  The planets \"tack\" and migrate outward when Saturn reaches the gap-opening mass and is caught in the 3:2 resonance with Jupiter.} {The aim is to test the viability of the Grand Tack model and  to study the dynamical evolution of Jupiter and Saturn during their growth from $10$ $\\mearth$ cores.} {We have performed numerical simulations using a grid--based hydrodynamical code. Most of our simulations assume an isothermal equation of state for the disk but a subset use a fully-radiative version of the code.} {For an isothermal disk the two phase migration of Jupiter and Saturn is very robust and independent of the mass-growth history of these planets provided the disk is cool enough.  For a radiative disk the we find some outcomes with two phase migrations and others with more complicated behavior.  We construct a simple, 1-D model of an evolving viscous disk to calculate the evolution of the disk's radiative properties: the disk transitions from radiative to isothermal from its outermost regions inward in time.} {We show that a two-phase migration is a natural outcome at late times even under the limiting assumption that isothermal conditions are required.  Thus, our simulations provide strong support for the Grand Tack scenario. } ",
        "introduction": "The current paradigm of the origin and evolution of the Solar System's giant planets follows several distinct stages: \\begin{enumerate} \\item The cores of Jupiter and Saturn form by accretion of planetesimals (e.g., Kokubo \\& Ida 1998; Levison et al. 2010).  The timescales of accretion are poorly constrained because as they grow the cores migrate due to both the back-reaction from planetesimal scattering (Fernandez \\& Ip 1984; Kirsh et al. 2009) and type I (tidal) interactions with the gaseous protoplanetary disk (Goldreich \\& Tremaine 1980; Ward 1986; Paardekooper et al. 2010). \\item Jupiter and Saturn's cores slowly accrete gas and each undergo a phase of rapid gas accretion (e.g., Mizuno 1980, Pollack et al. 1996).  The rapid phase of accretion is triggered when the mass in each planet's gaseous envelope is comparable to the core mass.  Runaway gas accretion lasts for roughly the local Kelvin-Helmholtz time and ceases when the planet opens an annular gap in the disk and transitions to type II migration (Lin \\& Papaloizou 1986; Ward 1997).  Because of its larger mass and smaller orbital radius, Jupiter is thought to have undergone runaway gas accretion before Saturn. \\item Once fully-formed, Saturn migrated faster (Masset \\& Papaloizou 2003), caught up to Jupiter, and was trapped in 3:2 resonance (Pierens \\& Nelson 2008).  Interestingly, this result is found to be a very robust outcome of the simulations, independent on the earlier evolution of Saturn's core. For instance, Pierens \\& Nelson (2008) investigated the scenario in which Saturn's core is initially trapped at the edge of Jupiter's gap and grows through  gas accretion from the disk. In that case, they demonstrated that although Saturn is temporarily locked in the 2:1 resonance with Jupiter, it becomes ultimately trapped in the 3:2 resonance. \\item Once Jupiter and Saturn are trapped in 3:2 resonance, the gaps carved by the two planets in the Solar Nebula overlap.  Saturn's gap is not as deep as Jupiter's (due to its smaller mass), and this causes Jupiter and Saturn to migrate outward while remaining in 3:2 resonance, provided that both the disk thickness and the disk viscosity are small enough (Masset \\& Snellgrove 2001; Morbidelli \\& Crida 2007).   Outward migration is stopped when the disk dissipates  or, if the disk is flared, at a critical distance where the disk is too thick and the structure of the two planets' common gap is compromised (Crida et al 2009).  \\footnote{We note that this step is somewhat uncertain because both the migration and accretion rates of giant planet cores should a roughly linear dependence on the planet mass.  Thus, we naively expect that Saturn's gas accretion should mimic Jupiter's as it migrates inward, meaning that Saturn should be roughly $1 M_J$ when it catches up to Jupiter, precluding outward migration, which requires a Saturn/Jupiter mass ratio of roughly 1/2 or smaller (Masset \\& Snellgrove 2001).  The solution to this problem is not clear: it may involve a change in the disk opacity to allow Saturn's rapid type III migration to last for longer than Jupiter's.  Such rapid migration has been invoked to explain the 3:2 resonant exoplanet system HD 45364 (Rein et al. 2010).  Of course it is reasonable to expect that this mechanism is probably not universal, since it depends on details such as the timing of core formation and the disk properties.  Thus, in many exoplanet systems \"Saturn\" would have reached $1 M_J$ and the two planets would not tacked and migrated outward.  The architecture of such systems naturally would not resemble the Solar System.  Understanding the statistical distribution of giant exoplanetary systems can therefore place constraints on their early evolution and the frequency of \"grand tacks\". } \\item After the dissipation of the gas disk, planetesimal-driven migration causes a large-scale spreading of the planets' orbits because Jupiter is the only planet for which the ejection of small bodies is more probable than inward scattering (Fernandez \\& Ip 1984; Hahn \\& Malhotra 1999).  In the \"Nice model\", Jupiter and Saturn are assumed to have formed interior to their mutual 2:1 resonance and, when they cross it, an instability is triggered that causes the Late Heavy Bombardment (Gomes et al. 2005).  Recent work has shown that the Nice model is still valid if more realistic initial conditions are used, with Jupiter and Saturn in 3:2 resonance and Uranus and Neptune also trapped in a resonant chain (Morbidelli et al. 2007; Batygin \\& Brown 2010).   The Nice model can reproduce the giant planets' final orbits (Tsiganis et al. 2005), the orbital distribution of Jupiter's Trojan asteroids (Morbidelli et al. 2005), and several other characteristics of the Solar System's small body populations. \\end{enumerate}  Although the detailed orbital evolution of the giant planets is not known, these steps explain their origin in broad strokes.  By assembling steps 2-4, Walsh et al. (2011) recently proposed a new model to explain the origin of the inner Solar System called the \"Grand Tack\".  In this model, Jupiter formed at $\\sim 2-5$ AU, migrated inward then \"tacked\" (i.e., changed the direction of its migration) at an orbital distance of $\\sim$1.5 AU when Saturn caught up and was trapped in 3:2 resonance and migrated back out past 5 AU.  Jupiter's tack at 1.5 AU truncates the inner disk of planetary embryos and planetesimals from which the terrestrial planets formed at about 1 AU.  This type of narrow truncated disk represents the only initial conditions known to satisfactorily reproduce the terrestrial planets, in particular the small mass of Mars compared with Earth (Wetherill 1978; Hansen 2009; Raymond et al. 2009; Walsh et al. 2011).  An additional success of the Grand Tack model is that the asteroid belt is naturally repopulated from two distinct populations corresponding to the C- and S- type asteroids.  At the end of the Grand Tack, the giant planets' orbits represent the initial conditions for the Nice model (Raymond et al., in prep.).  The goal of this paper is to test the viability of the evolution of Jupiter and Saturn in the Grand Tack model (Walsh et al 2011).  To accomplish this we use the GENESIS hydrocode to simulate the growth and migration of Jupiter and Saturn from $10$ $\\mearth$ cores.  With respect to previous simulations (Masset \\& Snellgrove 2001; Morbidelli \\& Crida 2007; Pierens \\& Nelson 2008), we consider a self-consistent scenario in which the cores of Jupiter and Saturn slowly grow to full-fledged gas giants by accreting gas from the disk.  We find that a two phase migration of Jupiter and Saturn is a very robust outcome in isothermal disks, but occurs in only one of two simulations in radiative disks.  We place our simulations in the context of the evolving Solar Nebula using a 1-D diffusion algorithm that differentiates between radiative and isothermal behavior.  Our results strongly favor a two-phase migration of Jupiter and Saturn, and support the Grand Tack.  The paper is organized as follows. In Sect. 2, we describe the hydrodynamical model. In Sect. 3, we present the results of isothermal simulations.  In Sect. 4 we present results of radiative simulations.  In Sect. 5 we construct a 1-D model of the Solar Nebula to show when the disk should be isothermal or radiative.  Finally, we discuss our results and draw conclusions in Sect. 6. ",
        "conclusions": "Our results indicate that Jupiter and Saturn probably underwent a two-phase, inward-then-outward migration.  In our simulations, Jupiter and Saturn start as $10$ $\\mearth$ cores and type I migrate; inward for isothermal disks, inward or outward for radiative disks.  In most cases the two cores become locked in 3:2 mean motion resonance (MMR).  At this point  or after a delay of 500-1000 orbits, we allowed Jupiter to start accreting gas from the disk.  When Jupiter reaches the gap-opening mass, it undergoes a phase of rapid inward migration as it clears out its gap (sometimes called type III migration; Masset \\& Papaloizou 2003) then settles into standard, type II migration.  Inward migration continues until Saturn accretes enough gas to reach the gap-opening mass itself.  At this point, Saturn's inward migration accelerates and is again trapped in the 3:2 MMR with Jupiter.  Outward migration of both giant planets is then triggered via the mechanism of Masset \\& Snellgrove (2001). Outward migration stops when either a) the disk dissipates (as in Sect. \\ref{sec:disp}), b) Saturn reaches the outer edge of the disk, or, c) if the disk is flared, the giant planets drop below the local gap-opening mass (e.g., Crida et al. 2009).  An additional stopping -- or at least slowing -- mechanism exists if the planets are unable to maintain a well-aligned resonant lock during migration.  For example, in simulation I1 Jupiter and Saturn's rate of outward migration slowed significantly when one resonant angle transitioned from libration to circulation (see Figs. 1 and 3).  Here, this appears to be due to a positive corotation torque exerted on Saturn by gas that polluted Jupiter and Saturn's common gap as the distance between the two planets increased during outward migration.  On longer timescales it is unclear if this mechanism would continue to slow down and eventually stop the outward migration.  In isothermal disks, the two phase migration of Jupiter and Saturn holds for almost the full range of parameters that we tested (Sect. 3).  The only situation for which this result does not hold is if the gaseous Solar Nebula is relatively thick ($h = H/r \\gtrsim 0.05$).  Both the disk's surface density profile and the value for the disk aspect ratio had an effect on the migration rate: disks with either shallower profiles or lower values of $h$ result in faster migration. Changing the disk's viscosity had little effect on the outcome, although we only tested a very small range.  In one simulation (I4; Sect. \\ref{sec:varyxj}), Saturn's core was pushed past the 2:1 MMR with Jupiter leading to a significant eccentricity increase for both planets before the resonance was crossed, Saturn was trapped in the 2:3 MMR and both planets migrated outward.  This dynamic phase of resonance crossing and eccentricity excitation is likely to be quite sensitive to the detailed properties of the disk (e.g., the scale height and viscosity) that determine the gap profile.  We performed two simulations in radiative disks with mixed outcomes (Sect. 4).  In the first case, Jupiter and Saturn started accreting together and so stayed relatively close to each other.  The planets became locked in the 3:2 MMR and migrated outward even faster than in isothermal simulations due to the small aspect ratio of the outer disk ($h \\approx 0.03$).  In the second case, Saturn started to accrete when Jupiter reached half its final mass (i.e., $x_J = 0.5$), by which time the two planets were beyond the 2:1 MMR.  The planets succeeded in breaking the 2:1 and 5:3 MMRs, became trapped in the 3:2 MMR and migrated outward only to undergo a dynamical instability putting the planets once again beyond the 2:1 MMR.  A more detailed study of outward migration in radiative disks is underway.  Using a simple 1-D model of an evolving Solar Nebula we showed that the disk should be optically thick at early times, then transition to optically thin from the outside-in during the late phases of its evolution.  At the orbital distance in question (1-10 AU), the disk transitions from radiative to isothermal behavior in the last 1-2 Myr of its evolution.  Thus, even if we make the ``pessimistic'' assumption that an isothermal disk is required for outward migration of Jupiter and Saturn, the disk fulfills the criteria for long-range outward migration in its late phases.  Outward migration of Jupiter and Saturn at this time is very likely provided the disk remains thin ($h \\lesssim 0.05$).   Our simulations therefore show that an inward-then-outward migration of Jupiter and Saturn is extremely likely, and that the last phase of outward migration probably coincided with the late phases of the dissipation of the Solar Nebula.  This is of particular interest because the two phase migration of Jupiter and Saturn helps resolve a long-standing problem in terrestrial planet formation.  For over 20 years, simulations of terrestrial accretion have been unable to reproduce Mars' relatively small mass ($0.11$ $\\mearth$; Wetherill 1978, 1991; Chambers 2001; Raymond et al. 2009).  This problem arises because, in a Solar Nebula that varies smoothly in orbital radius, there is a comparable or larger amount of mass in the vicinity of Mars than the Earth.  For Mars to be so much smaller than Earth, most of the mass between roughly 1-3 AU must be removed (e.g., Raymond et al. 2006, 2009; O'Brien et al. 2006).  Several mechanisms have been proposed to remove this mass, including strong secular resonances (Thommes et al. 2008, Raymond et al. 2009) and a narrow dip in the surface density caused by a radial dependence of the disk's viscosity (i.e., a dead zone; Jin et al. 2008).  However, the problem is most easily and much better solved if the terrestrial planets did not form from a wide disk of planetary embryos but instead from a narrow annulus extending only from 0.7-1 AU (Wetherill 1978; Chambers 2001; Hansen 2009).  In that case, Mars' small mass is simply an edge effect: Mars is small was built from  one or perhaps a few embryos that were scattered beyond the edge of the embryo disk (this is also the case for Mercury, which was scattered inward beyond the inner edge of the embryo disk).  In contrast, Earth and Venus formed within the annulus and are consequently much more massive.  Simulations of terrestrial planet formation can quantitatively reproduce the orbits and masses of all four terrestrial planets as well as their radial distribution (Hansen 2009).  The flaw in simulations of terrestrial planet formation in truncated disks is that they had no justification for the truncation; the ad-hoc initial conditions were simply chosen because they provided a good fit to the actual terrestrial planets (Hansen 2009).  The two phase migration of Jupiter and Saturn provides such a justification via the Grand Tack model of Walsh et al. (2011).  If Jupiter's turnaround point was at $\\sim 1.5$ AU then it would have naturally truncated the inner disk of embryos and planetesimals at about 1 AU -- in most of our simulations Jupiter indeed tacked at roughly this distance.  As expected, the terrestrial planets that form from this disk quantitatively reproduce the actual terrestrial planets (Walsh et al. 2011).  The Grand Tack model also provides the best explanation to date for the observed dichotomy between the inner and outer asteroid belt (Gradie \\& Tedesco 1982).  Thus, the present-day Solar System appears to bear the imprint of a two phase migration of Jupiter and Saturn.   Our hydrodynamical simulations provide support for the Grand Tack scenario.  As with any numerical study, our simulations do not fully represent reality.  The aspect of our simulations that is probably the least realistic is the gas accretion onto the giant planets' cores. In our simulations, gas accretion onto Jupiter and Saturn is extremely fast.  Once accretion starts, Jupiter and Saturn reach their final masses in only a few thousand years, whereas the Kelvin-Helmholtz time in protoplanetary disks is more like $\\sim 10^5$ years.  In addition, accretion onto growing giant planet cores requires transferring gas through circum-planetary accretion disks whose physical properties are poorly constrained (e.g., Ward \\& Canup 2010).  Once the planets reached their actual masses we artificially turned off gas accretion. If, during the outward migration Saturn accreted enough gas to carve a gap as deep as Jupiter's then Saturn's outer lindblad torque would balance Jupiter's inner lindblad torque, outward migration would stop and the planets would turn back around and migrate inward.  The impact of a more realistic accretion history on Jupiter and Saturn's migration remains an open question, in particular with regards to the interplay between gas accretion and the dispersal of both the cimcumstellar and cicumplanetary disks."
    },
    "1107/1107.0842_arXiv.txt": {
        "abstract": "{} {We investigate the  structure and shape of the photospheric and molecular layers of the atmospheres of four Mira variables.} {We obtained near-infrared $K$-band spectro-interferometric observations of the Mira variables \\object{R~Cnc}, \\object{X~Hya}, \\object{W~Vel}, and \\object{RW~Vel} with a spectral resolution of about 1500 using the AMBER instrument at the VLTI. We obtained concurrent JHKL photometry using the the Mk~II instrument at the SAAO. } {The Mira stars in our sample are found to have wavelength-dependent visibility values that are consistent with earlier low-resolution AMBER observations of S Ori and with the predictions of dynamic model atmosphere series based on self-excited pulsation models. The corresponding wavelength-dependent uniform disk (UD) diameters show a minimum near the near-continuum bandpass at 2.25\\,$\\mu$m. They then increase by up to 30\\% toward the H$_2$O band at 2.0\\,$\\mu$m and by up to 70\\% at the CO bandheads between 2.29\\,$\\mu$m and 2.48\\,$\\mu$m. The dynamic model atmosphere series show a consistent wavelength-dependence, and their parameters such as the visual phase, effective temperature, and distances are consistent with independent estimates. The closure phases have significantly wavelength-dependent and non-zero values at all wavelengths indicating deviations from point symmetry. For example, the R Cnc closure phase is 110\\degr\\ $\\pm$ 4\\degr\\ in the 2.0\\,$\\mu$m H$_2$O band, corresponding for instance to an additional unresolved spot contributing 3\\% of the total flux at a separation of $\\sim$ 4\\,mas. } {Our observations are consistent with the predictions of the latest dynamic model atmosphere series based on self-excited pulsation models. The wavelength-dependent radius variations are interpreted as the effect of molecular layers lying above the photosphere. The wavelength-dependent closure phase values are indicative of deviations from point symmetry at all wavelengths, thus a complex non-spherical stratification of the extended atmosphere. In particular, the significant deviation from point symmetry in the H$_2$O band is interpreted as a signature on large scales (there being a few across the stellar disk) of inhomogeneities or clumps in the water vapor layer. The observed inhomogeneities might possibly be caused by pulsation- and shock-induced chaotic motion in the extended atmosphere. } ",
        "introduction": "Mira stars are long-period, large-amplitude variable stars on the asymptotic giant branch (AGB). Mass loss becomes increasingly important toward the tip of the AGB, before the star evolves to the planetary nebula (PN) phase, where a great diversity of morphologies is seen. Contemporary astrophysical questions related to the study of Mira variables include the pursuit of the detailed mass-loss mechanism on the AGB, including the effects of pulsation and shock fronts on the structure and morphology of the extended atmosphere, and of the mechanism shaping the observed PN morphologies. Near-infrared interferometry has traditionally been used to characterize AGB star atmospheres. In particular, observations using the IOTA interferometer have uncovered the wavelength-dependence of Mira star diameters using a few bandpasses with spectral resolutions of up to $\\lambda/\\Delta\\lambda\\sim25$, providing observational evidence of molecular layers lying outside the photospheric layers (e.g.; Mennesson et al.  \\cite{mennesson02}; Perrin et al. \\cite{perrin04}). These molecular layers have also been present in theoretical dynamic model atmospheres (Hofmann et al. \\cite{hofmann98}, Ireland et al. \\cite{ireland04a,ireland04b}). $H$-band interferometry at the IOTA interferometer in the broad band (Ragland et al. \\cite{ragland06}), as well as in three filters with $\\lambda/\\Delta\\lambda\\sim15$ (Ragland et al. \\cite{ragland08}, Pluzhnik et al. \\cite{pluzhnik09}) have revealed non-zero closure phases for several Mira variables, which reflect asymmetric brightness distributions of the photosphere and/or the envelope around the star. Wittkowski et al. (\\cite{wittkowski08}) presented the first VLTI/AMBER near-infrared spectro-interferometric observation of an AGB star, providing continuous wavelength coverage from 1.29\\,$\\mu$m to 2.32\\,$\\mu$m with a spectral resolution of $\\lambda/\\Delta\\lambda\\sim35$. The data showed visibility and diameter variations as a function of wavelength that generally confirmed the predictions of dynamic model atmospheres, where the diameter variations can be understood as the effects from atmospheric molecular layers (most importantly H$_2$O, CO, and TiO). Here, we present the first VLTI/AMBER observations of Mira variables using its $K$ medium resolution modes with a spectral resolution of $\\lambda/\\Delta\\lambda\\sim1500$, and a comparison to the newly available {\\tt CODEX} dynamic model-atmosphere series by Ireland et al. (\\cite{ireland08,ireland11}), which are based on self-excited pulsation models. ",
        "conclusions": "The four Mira variables of our sample exhibit consistent characteristic wavelength dependences of the visibility and consequently the corresponding uniform disk diameter that are consistent with those of earlier low-resolution AMBER data of the Mira variable S~Ori and the predictions of the {\\tt P/M} and {\\tt CODEX} dynamic model atmosphere series. Here, the newly available {\\tt CODEX} series provides a closer agreement with the data than the earlier {\\tt P/M} series. This result confirms that the wavelength-dependent angular diameter is caused by the atmospheric molecular layers, here most importantly H$_2$O and CO, as they are naturally included in the dynamic model atmosphere series. Concurrent $JHKL$ photometry obtained at the SAAO was used to derive $T_\\mathrm{eff}$ based on the integrated bolometric flux and the fitted Rosseland-mean angular diameter. Parameters of the best-fit model atmospheres, such as visual phase, effective temperature, and distances are consistent with independent estimates, which provides additional confidence in the {\\tt CODEX} modeling approach.  The closure phase functions of our targets exhibit non-zero values at all wavelengths with a wavelength dependence that also correlates with the positions of the H$_2$O and CO bands. This result indicates a complex non-spherical stratification of the extended atmosphere of Mira variables. The most significant deviation from point symmetry is observed in the H$_2$O band around 2.0\\,$\\mu$m, in particular for the clearly resolved target R~Cnc. We interpret this signal as an indication of inhomogeneities or clumps within the water vapor layer on scales of a few resolution elements across the stellar disk. These inhomogeneous water shells have also been detected for the Mira variable U Ori (Pluzhnik et al. \\cite{pluzhnik09}) and the symbiotic Mira variable R Aqr (Ragland et al. \\cite{ragland08}). The deviations from point symmetry at the near-continuum bandpass may be related either to photospheric convection cells (cf. Freytag \\& H\\\"ofner \\cite{freytag08}) or the inhomogeneities of molecular layers that may also contaminate this bandpass.  Inhomogeneous or clumpy molecular layers may be expected as the result of chaotic motion induced by the interaction of pulsation and shock fronts with the extended atmosphere. The CODEX models indicate that the outer mass zones outward of 1.5--2 Rosseland radii are only loosely connected to the stellar pulsation (cf. Fig. 1 of Ireland et al. \\cite{ireland08}). Likewise, outer mass zones on different sides of the star may be only weakly correlated with each other and may have different extensions. Icke et al. (\\cite{icke92}) described as well that the outer layers of an evolved AGB star may respond with chaotic motion to the pulsations that originate in the stellar interior. Observations of water vapor layers may be particularly sensitive to these effects, but other molecules, such as CO, that are expected to be plentiful in the shocked region of the atmosphere (cf. Cherchneff \\cite{cherchneff06}) would also be affected by this large-scale chaotic motion. The resulting clumpy structure of molecular layers may explain our complex wavelength-dependent closure phase signal.  Further interferometric campaigns with high spatial and spectral resolution are clearly needed to characterize in detail the morphology of atmospheric molecular layers in Mira variables."
    },
    "1107/1107.0910_arXiv.txt": {
        "abstract": "Very high energy (VHE; $>$100 GeV) $\\gamma$-rays are expected to be emitted from the vicinity of super-massive black holes (SMBH), irrespective of their activity state. In the magnetosphere of rotating SMBH, efficient acceleration of charged particles can take place through various processes. These particles could reach energies up to $E\\sim10^{19}$eV. VHE \\gr emission from these particles is then feasible via leptonic or hadronic processes. Therefore passive systems, where the lack of a strong photon field allows the VHE \\grs to escape, are expected to be detected by Cherenkov telescopes. We present results from recent VHE experiments on the passive SMBH in the nearby elliptical galaxy NGC 1399. No \\gr signal has been found, neither by the \\hess experiment nor in the Fermi data analyzed here. We discuss possible implications for the physical characteristics of the system. We conclude that in a scenario where particles are accelerated in vacuum gaps in the magnetosphere, only a fraction $\\sim 0.3$ of the gap is available for particle acceleration, indicating that the system is unlikely to be able to accelerate protons up to $E\\sim10^{19}$eV. ",
        "introduction": "Spheroidal systems (such as elliptical galaxies, lenticular galaxies, and early-type spiral galaxies with massive bulges) are commonly believed to host super-massive black holes with masses in the range $\\mbh= (10^6 - 10^9)~M_\\odot$ in their central region \\citep[e.g.,][]{rich,ferr-ford}. Many SMBH are {\\it active} and give rise to emission over a wide range of frequencies. Some have been detected from radio to high energy \\grs ($\\sim1$ GeV). The SMBH is essential for the production and stability of relativistic jets. In many objects the inner, compact jets emit VHE \\gr emission through various leptonic and/or hadronic processes. In blazar-type objects, where the jet is pointing towards the observer, VHE \\gr emission is frequently observed \\citep[e.g.,][]{hin-hof,weekes}. Often, this can be successfully modeled by non-thermal electrons upscattering either their own synchrotron radiation or upscattering an external photon field, although a hadronic origin \\citep[e.g.,][]{aharonian00, muc-prot} cannot be excluded. In these cases the gamma-ray emission is released through the jets which in turn are powered by the SMBH.  VHE emission may also be indirectly related to SMBH  independent of any electromagnetic activity. The presence of a SMBH steepens the potential well and hence the density profile of dark matter in the central regions of galaxies. The rate of annihilation of Dark Matter particles \\citep[e.g.][]{darkmatter} will thus be increased resulting in enhanced emission of the gamma-rays generated in the annihilation.  Several models for the direct production of \\gr emission in the vicinity of SMBH have been proposed \\citep[see e.g.,][]{cascadeinagn,mastichiadis,slane,boldt-gosh,lev,neron,neron-aha,rieger,istomin-sol,osm10,lev11}. Possibly, some of these mechanisms could also be responsible for acceleration of cosmic-rays to energies $E\\sim10^{19}$ eV and beyond. In order to avoid attenuation on circumnuclear fields, VHE photons could escape only if the SMBH does not produce too much low-energy radiation. The SMBH would have to be passive at low energies, i.e. most of the radiative losses would have to occur at high energies. In all cases a large mass of the central object is an important characteristic for generating a high VHE flux.  Correlations involving the SMBH masses and properties of their host galaxies have been investigated by many authors. In particular $\\mbh$ is found to be linked to the central stellar velocity dispersion \\citep[e.g.,][]{gebhardt} or to the mass of the host galaxy bulge \\citep[e.g.,][]{magorrian}.  These observational scaling laws, in addition to confirming the ubiquity of SMBH, have suggested that their activity  is tightly linked to the evolution of their host galaxies \\citep[e.g.][]{ferr-ford}.  During the early stages of galaxy evolution SMBH accrete matter at high rates and are observed as bright QSO (Quasi-Stellar Objects). Even if such systems generate VHE gamma-rays, the dense photon fields associated with the Quasar phase would pair-absorb the VHE radiation. The average radiative output at low photon energies (e.g., in the optical band) decays from redshift $z>3$ to $z=0$ by almost 2 orders of magnitude. The majority of SMBH in the local universe are hosted in systems of low accretion rate and are therefore not embedded in dense radiation fields. In {\\it passive} systems (i.e., SMBH hosted in nuclei without bright signatures of broad-band activity and very low luminosity in longer wavelengths). If VHE gamma-rays, if generated, can escape from the nuclear region without suffering from strong absorption via photon-photon pair absorption.  While the detection of VHE gamma-rays in blazar-type systems is facilitated by the superluminal motion (the apparent luminosity is boosted and the optical depth related to photon-photon absorption is reduced, see \\citet{blaz_sch}), this is not the case for non-blazar systems. Hence, proximity and low luminosity in the IR/optical domain increase the possibility of a detection in the VHE band. Observations of passive systems can hence contribute to our understanding of the physics and properties of galactic nuclei. In addition they might give us an insight on Ultra High Energy Cosmic Ray (UHECR; $E>4 \\times 10^{18}$ eV) sources.  Here we present GeV limits and discuss observations of the passive SMBH in the core of NGC 1399, the central galaxy of the Fornax cluster, that were conducted with the \\hess Cherenkov telescope array \\citep{mio_icrc} and discuss its implications for gap-type particle acceleration and emission models. ",
        "conclusions": "If efficient gap-type particle acceleration and curvature emission occurs in NGC 1399, VHE $\\gamma$-ray emission should have been detected by Fermi and/or \\hess According to eq.~(\\ref{eq:luminosity_adaf}), the non-detection of NGC 1399 thus either suggests that (i) the strength of the ordered, magnetospheric field component is only a small fraction ($\\leq 0.1$) of the disk magnetic field value estimated above, (ii) on average only a small fraction of the gap ($[h/r_g]^2\\ll1$) is available for particle acceleration and/or (iii) the charge density $n_e$ in the vicinity of the black hole is much smaller than the Goldreich-Julian density, i.e. $(n_e/n_{\\rm GJ})\\ll1$.  Option (i) would be incompatible with the jet power $L_{\\rm jet} \\simeq 10^{42}$ erg/s \\citep{shurkin} being provided by a Blandford-Znajek-type process. Option (ii) would require that the charge density around the black hole exceeds the Goldreich-Julian density $n_{\\rm GJ}\\simeq 2 \\times 10^{-3} (B/10^3\\mathrm{G})$ cm$^{-3}$. This could be the case if efficient pair production occurs near to the black hole \\citep[e.g.,][]{moc11,lev11}. The accretion rate for NGC 1399 inferred above \\citep[cf. also][]{narayan02} seems indeed close to the critical value $\\dot{m}_c \\simeq 2 \\times 10^{-4}$ where annihilation of MeV photons in a two-temperature ADAF could lead to the injection of seed charges with density $n_e \\geq n_{\\rm GJ}$ \\citep{lev11}. If the accretion rate would be sufficiently high, i.e., $\\dot{m} >\\dot{m}_c$ ensuring $n_e > n_{\\rm GJ}$, a substantial part of the gap is expected to be screened, suggesting $(h/r_g)^2 \\ll 1$. This could make the anticipated VHE output, eq.~(\\ref{eq:luminosity_adaf}), consistent with the VHE upper limits derived above. However, it would also imply that the available potential, eq.~(\\ref{V_effective}), is reduced down to a level where UHE proton acceleration will no longer be possible. On the other hand, if the accretion would be such that $\\dot{m} < \\dot{m}_c$, implying $n_e < n_{\\rm GJ}$ (option [iii]), fully-developed gaps ($h\\sim r_g$) may exist. This could again make $L_{\\gamma}$ consistent with the observationally inferred VHE upper limits. However, an additional plasma source would then be needed to establish the force-free outflow believed to be present on larger scales. One plausible scenario relates to pair cascade formation in charge-starved magnetospheric regions due to the absorption of inverse Compton up-scattered photons \\citep{lev11}. When applied to NGC 1399, the estimated multiplicity appears indeed sufficiently high ($M \\simgt 10^3$) to allow the pair density to approach the Goldreich-Julian density. As the optical depth for pair production across the gap is larger than unity even if the gap is not fully restored (i.e., for $h<r_g$), cascade formation can occur on scales $< r_g$ and thereby limit the potential gap size. Detailed modelling would be required to self-consistently calculate the gap size and associated VHE emission. On the other hand, the estimated VHE \\gr output $L_{\\gamma}$ is roughly a fraction $(h/r_g)^2$ of the maximum Blandford-Znajek jet power $L_\\mathrm{BZ}$. Assuming $L_{\\rm BZ}=L_{\\rm jet}$, the VHE flux upper limits imply that $(h/r_g)^2 \\simlt 0.1$, again yielding conditions not conducive for efficient proton acceleration beyond $4 \\times 10^{18}$ eV.  Nearby, passive supermassive black holes are often believed to be prime candidates for magnetospheric, gap-type particle acceleration and emission scenarios. The non-detection of VHE gamma-ray emission in NGC 1399 by current VHE instruments now suggests that this object is unlikely to be an efficient UHECR proton accelerator."
    },
    "1107/1107.2638_arXiv.txt": {
        "abstract": "Monte Carlo simulations of neutron-rich matter of relevance to the inner neutron-star crust are performed for a system of $A\\!=\\!5,000$ nucleons. To determine the proton fraction in the inner crust, numerical simulations are carried out for a variety of densities and proton fractions. We conclude---as others have before us using different techniques---that the proton fraction in the inner stellar crust is very small. Given that the purported {\\sl ``nuclear pasta''} phase in stellar crusts develops as a consequence of the long-range Coulomb interaction among protons, we question whether pasta formation is possible in such proton-poor environments. To answer this question, we search for physical observables sensitive to the transition between spherical nuclei and exotic pasta structures. Of particular relevance is the static structure factor $S(k)$---an observable sensitive to density fluctuations. However, no dramatic behavior was observed in $S(k)$. We regard the identification of physical observables sensitive to the existence---or lack-thereof---of a pasta phase in proton-poor environments as an open problem of critical importance. ",
        "introduction": "\\label{Introduction}  Neutron stars are unique laboratories for the study of matter under extreme conditions of density and isospin asymmetry~\\cite{Weber:1999,Glendenning:2000}. Indeed, the conditions in the interior of neutron stars are so extreme that they are unattainable in terrestrial laboratories. Thus, neutron stars---and the exotic phases within---owe their existence to the presence of enormous gravitational fields. To maintain hydrostatic equilibrium throughout the star, these enormous gravitational fields must be balanced by the pressure support of its underlying constituents. This results in an enormous dynamic range of pressures and densities that enables one to probe the equation of state (EOS) far away from the equilibrium density of normal nuclei~\\cite{Piekarewicz:2009zz}.  Neutron stars contain a non-uniform crust above the uniform liquid mantle (or stellar core). The core is comprised of a uniform assembly of neutrons, protons, electrons, and muons packed to densities that may exceed that of normal nuclei by up to an order of magnitude.  The highest density attained in the core depends critically on the equation of state of neutron-rich matter which at those high densities is presently poorly constrained. The core accounts for almost all of the mass and most of the size of the neutron star. However, at densities of about half of normal nuclear density, the uniform core becomes unstable against cluster formation. At these ``low'' densities the average inter-nucleon separation increases to such an extent that it becomes energetically favorable for the system to segregate into regions of normal density (nuclear clusters) and regions of low density (dilute neutron vapor).  Such a clustering instability signals the transition from the uniform liquid core to the non-uniform crust.  Note, however, that the precise value of the crust-to-core transition density is presently unknown, as it is sensitive to the poorly constrained density dependence of the symmetry energy~\\cite{Horowitz:2000xj}.  The solid crust is divided into an outer and an inner region. In particular, the outer crust spans a region of about seven orders of magnitude in density (from about $10^{4}{\\rm g/cm^{3}}$ to $4\\times 10^{11}{\\rm g/cm^{3}}$~\\cite{Baym:1971pw,Ruester:2005fm,RocaMaza:2008ja}). Structurally, the outer crust is comprised of a Coulomb lattice of neutron-rich nuclei embedded in a uniform electron gas. As the density increases---and given that the electronic Fermi energy increases rapidly with density---it becomes energetically favorable for electrons to capture into protons. This results in the formation of Coulomb crystals of progressively more neutron-rich nuclei. Eventually, the neutron-proton asymmetry becomes too large for the nuclei to absorb any more neutrons and the excess neutrons go into the formation of a dilute neutron vapor; this signals the transition from the outer to the inner crust.  In this contribution we are interested in modeling the structure of the inner crust.  The outer-to-inner crust transition density is predicted to occur at about $4\\times 10^{11}{\\rm g/cm^{3}}$~\\cite{Baym:1971pw,Ruester:2005fm,RocaMaza:2008ja}. At this---{\\sl neutron-drip}---density the neutron-rich nucleus (${}^{118}$Kr) that comprises the crystalline lattice is unable to retain any more neutrons. Thus, the top layer of the inner crust consists of a Coulomb crystal of neutron-rich nuclei immersed in a uniform electron gas and a dilute---likely superfluid---neutron vapor. In contrast, at the bottom layer of the inner crust the density has become high enough (of the order of $10^{14}{\\rm g/cm^{3}}$) to {\\sl ``melt''} the crystal into a uniform mixture of neutrons, protons, and electrons. Yet the transition from the highly-ordered crystal to the uniform liquid is both interesting and complex. This is because distance scales that were well separated in both the crystalline phase (where the long-range Coulomb interaction dominates) and in the uniform phase (where the short-range strong interaction dominates) become comparable. This unique situation gives rise to {\\sl ``Coulomb frustration''}.  Frustration, a phenomenon characterized by the existence of a very large number of low-energy configurations, emerges from the impossibility to simultaneously minimize all {\\sl ``elementary''} interactions in the system. Whereas protons are correlated at short distances by attractive strong interactions, they are anti-correlated at large distances because of the Coulomb repulsion. Whenever these short and large distance scales are well separated (as in the outer crust) protons bind into nuclei that are then segregated in a crystal lattice. However, as these length scales become comparable---at densities of about $10^{13}\\!-\\!10^{14}$~g/cm$^3$ as in the inner crust---competition among the elementary interactions results in the formation of complex topological structures collectively referred to as {\\sl ``nuclear pasta''}. Given that these complex structures are very close in energy, it has been speculated that the transition from the highly ordered crystal of spherical nuclei to the uniform phase must proceed through a series of changes in the dimensionality and topology of these structures~\\cite{Ravenhall:1983uh,Hashimoto:1984}. Moreover, due to the preponderance of low-energy states, frustrated systems display an interesting and unique low-energy dynamics.  Interestingly enough, a seemingly unrelated condensed-matter problem---{\\sl the strongly-correlated electron gas}---appears to be connected to the nuclear pasta. In the case of the electron gas, one aims to characterize the transition from the low-density Wigner crystal (where the Coulomb potential dominates) to the uniform high-density Fermi liquid (where the kinetic energy dominates)~\\cite{Fetter:1971}.  It has been argued that instead of the expected first order phase transition, the transition is mediated by the emergence of {\\sl ``microemulsions''}, namely, exotic (``pasta-like'') structures.  Indeed, it has been shown that in {\\sl two-spatial} dimensions first-order phase transitions are forbidden in the presence of long-range ({\\sl e.g.,} Coulomb) forces~\\cite{Jamei:2005}. One should note that no generalization of this theorem to three dimensions exists. Hence, although the existence of pasta phases in both core-collapse supernovae and neutron stars may be plausible~\\cite{Ravenhall:1983uh,Hashimoto:1984}, there is no guarantee that they exist. Actually, in a recent publication Oyamatsu and Iida have shown that pasta formation may not be universal and that its existence (or lack-thereof) is intimately related to the density dependence of the symmetry energy~\\cite{Oyamatsu:2006vd}.  In particular, the authors concluded that pasta formation requires models with a soft symmetry energy, namely, ones that increase slowly with density.  Although the complex dynamics of the electron gas may shed light on the possible emergence of pasta phases in the stellar crust, an essential difference between these two systems remains. Whereas the constituents of the electron gas ({\\sl i.e.,} electrons) are all electrically charged and thus experience long-range forces, the neutrons in the crust interact solely via short-range forces.  Thus, a large neutron fraction in the inner crust could hinder the formation of the nuclear pasta. Indeed, pure neutron matter does not cluster. So given that the proton fraction in the crust is highly sensitive to the density dependence of the symmetry energy, the emergence of pasta phases involves a delicate interplay between: (i) the symmetry energy (which controls the proton fraction), (ii) the surface tension (which favors spherical nuclei), and (iii) the Coulomb interaction (which favors deformation). It is this interesting and unique problem that we propose to study here via numerical simulations.  Monte-Carlo simulations will be performed directly in terms of the nucleon constituents. Although we have already performed simulations of this kind~\\cite{Horowitz:2004yf,Horowitz:2004pv,Horowitz:2005zb}, in the present manuscript we aim to improve them in two critical areas.  First, in our previous work a screened Coulomb interaction between the protons was assumed with a screening length fixed at $10$~fm.  Although it was argued that no major qualitative changes are expected from such an approximation, in this contribution we evaluate the Coulomb interaction exactly via an Ewald summation.  For completeness, an appendix has been included where a detailed derivation of the Ewald summation is presented for a system of protons embedded in a uniform electron gas. Second, our earlier simulations were all performed at the single proton fraction of $Y_{p}\\!=\\!0.2$. This value was chosen as it is intermediate between the large values attained in core-collapse supernovae and the low proton fractions characteristic of the inner crust of cold---fully catalyzed---neutron stars.  Instead, in this work we will simulate neutron-rich matter for a variety of densities and proton fractions, albeit for a relatively small number of nucleons ($A\\!=\\!5,000$). In this way, the optimum proton fraction will be determined by imposing $\\beta$-equilibrium.  We should note that for the large proton fractions characteristic of core-collapse supernovae ($Y_{p}\\!=\\!0.3$--$0.5$) there appears to be agreement that the transition from the ordered Coulomb crystal to the uniform phase must proceed via intermediate pasta phases (such as rods, slabs, tubes, {\\sl etc.}). This agreement has been reached whether one uses numerical simulations~\\cite{Horowitz:2004yf,Horowitz:2004pv, Horowitz:2005zb,Watanabe:2003xu,Watanabe:2004tr, Watanabe:2009vi} or mean-field approximations with a variety of effective interactions~\\cite{Maruyama:2005vb,Avancini:2008zz,Avancini:2008kg, Newton:2009zz,Shen:2011kr}.  What is unclear, however, is whether such exotic pasta shapes will still develop in the proton-poor environment of the inner stellar crust. For example, mean-field models that impose $\\beta$-equilibrium predict proton fractions at densities of relevance to the inner crust of only a few percent~\\cite{Maruyama:2005vb,Avancini:2008zz,Avancini:2008kg, Shen:2011kr}. To our knowledge, no numerical simulations have been carried out with proton fractions less than $Y_{p}\\!=\\!0.1$~\\cite{Watanabe:2003xu}. Thus, our manuscript revolves around the following three fundamental questions: (a) Do numerical simulations support the very low proton fractions predicted by mean-field calculations? (b) Does the transition from a highly-ordered crystal of spherical nuclei to a uniform Fermi liquid must proceed via intermediate pasta phases even when the proton fraction is very small? (c) Can one identify a set of robust physical observables that are sensitive to the formation of the nuclear pasta?  We have organized the manuscript as follows. In Sec.~\\ref{Formalism} we introduce the model and develop the formalism necessary to carry out the Monte-Carlo simulations. Results are presented in Sec.~\\ref{Results} for the potential energy as a function of density and proton fraction. The optimal proton fraction is then determined by minimizing the total energy per nucleon of the system.  The pair-correlation function (for both protons and neutrons) and the corresponding static-structure factor are computed with the aim of identifying sensitive observables to the formation of the nuclear pasta. Our conclusions and suggestions for future work are presented in Sec.~\\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  Uniform neutron-rich matter at sub-saturation density is unstable against cluster formation. At densities below the neutron-drip line the ground-state of neutron-rich matter consists of a Coulomb crystal of spherical (neutron-rich) nuclei embedded in a uniform sea of electrons. It has been speculated that the transition from the highly-ordered crystal to the uniform Fermi liquid is mediated by an exotic phase known as {\\sl ``nuclear pasta''}. The nuclear-pasta phase---believed to exist in the inner crust of neutron stars---is characterized by the emergence of exotic nuclear shapes of various topologies that coexist with a dilute vapor of neutrons and electrons. Critical to the formation of the pasta is the long-range Coulomb interaction among protons, as it promotes the deformation of the clusters and the development of exotic topological structures (see Fig.~\\ref{Fig3}).  Given this essential fact, it is important to establish whether the proton fraction in the inner stellar crust is large enough to favor pasta formation.  To compute the proton fraction in the inner stellar crust we relied on Monte Carlo simulations assuming that nucleons interact via a two-body, short-range nuclear interaction and a long-range Coulomb repulsion. We employed a simple $NN$ interaction identical to the one introduced in Ref.~\\cite{Horowitz:2004yf} that was (approximately) fitted to a few ground-state properties of finite nuclei, the saturation properties of symmetric nuclear matter, and that precludes the binding of pure neutron matter. Yet we improved on the work of Ref.~\\cite{Horowitz:2004yf} by treating the long-range Coulomb interaction exactly via an Ewald summation.  Monte Carlo simulations were performed for a system of $A\\!=\\!5,000$ nucleons at a variety of densities and proton fractions (note that the simulations in Ref.~\\cite{Horowitz:2004yf} were limited to the single proton fraction of $Y_{p}\\!=\\!0.2$). By doing so, the optimum proton fraction (at each density) was determined by imposing $\\beta$-equilibrium. The optimum proton fraction emerges from a delicate competition between the nuclear symmetry energy---which favors proton fraction of $Y_{p}\\!\\lesssim\\!1/2$---and the Fermi energy of the electrons---which favors $Y_{p}\\!=\\!0$. For the model employed here, the rapid increase of the electronic contribution with proton fraction dominates over the nuclear contribution. This leads to {\\sl very small proton fractions at all densities of relevance to the inner crust}. For example, at a density of $\\rho\\!=\\!0.03~{\\rm fm}^{-3}$ the proton fraction is $Y_{p}\\!\\lesssim\\!0.02$---at least a factor of 10 smaller than assumed in Ref.~\\cite{Horowitz:2004yf}. Such small values for the proton fraction are consistent with those obtained using vastly different approaches (see, for example, Refs.~\\cite{Maruyama:2005vb,Avancini:2008zz,Avancini:2008kg,Shen:2011kr} and references therein). Ultimately, the small proton fractions obtained in most approaches appears to be a direct consequence of the rapid increase with density of the electronic contribution.  Given these results and the predominant role played by the Coulomb interaction in the formation of the nuclear pasta, it is only natural to ask whether such exotic structures can develop in the proton-poor environment of the inner stellar crust. To answer this question it becomes critical to identify observables that may be sensitive to pasta formation. Particularly useful in this case are observables sensitive to fluctuations, as these tend to increase rapidly near phase transitions. In this contribution we have focused on the static structure factor $S(k)$ which is sensitive to density fluctuations.  Moreover, the static structure factor is interesting as it provides a natural meeting place for theory, experiment, and computer simulations. Theoretically, $S(k)$ is the Fourier transform of the pair-correlation function $g(r)$---a quantity that is amenable to computer simulations. Experimentally, the proton (neutron) static structure factor can actually be measured in electron (neutrino) scattering experiments. In this work we have computed the pair-correlation function---for both protons and neutrons---as a function of proton fraction for a fixed density of $\\rho\\!=\\!0.03~{\\rm fm}^{-3}$.  In turn, the static structure factors were obtained by performing the required Fourier transforms. Although both $g(r)$ and $S(k)$ display interesting behavior that may be indicative of significant structural changes in the system, we find no clear evidence either in favor or against the formation of the nuclear pasta.  In particular, the static structure factor at zero momentum transfer $S(k\\!=\\!0)$ does not display the large enhancement characteristic of a phase transition. Perhaps what is required to clearly identify the onset and evolution of the pasta phase are dynamical observables, such as various transport coefficients.  To do so, one must rely on Molecular Dynamics (rather than Monte Carlo) simulations. A calculation along these lines was reported in Ref.~\\cite{Horowitz:2008vf} where both the shear viscosity and thermal conductivity were computed. Perhaps surprisingly, no dramatic increase in the viscosity was observed as the system evolves from spherical to exotic nuclear shapes. Clearly, more effort should---and will---be devoted along these lines.  In conclusion, our work revolved around three fundamental questions. First, what are typical values for the proton fraction in the inner stellar crust? Second, can exotic pasta structures develop in a proton-deficient environment? Finally, which physical observables are particularly sensitive to pasta formation?  We have found---as others have before us using vastly different approaches---that the proton fraction in the inner crust is very small indeed. However, at this moment it is unclear whether exotic pasta structure can develop in such proton-poor environments. In order to answer this question one would have to identify physical observables that are sensitive to the formation of the pasta. This task, however, has so far proven elusive. In an effort to answer these open questions we plan to perform Molecular Dynamics simulations with a very large number of particles. By doing so we will be able to calculate both static and dynamic observables. We trust that hidden among these host of observables will be at least one that will respond dramatically to the formation of the nuclear pasta."
    },
    "1107/1107.0317_arXiv.txt": {
        "abstract": "We analyze the dependence of galaxy structure (size and Sersic index) and mode of star formation ($\\Sigma_{SFR}$ and $SFR_{IR}/SFR_{UV}$) on the position of galaxies in the SFR versus Mass diagram.  Our sample comprises roughly 640000 galaxies at $z \\sim 0.1$, 130000 galaxies at $z \\sim 1$, and 36000 galaxies at $z \\sim 2$.  Structural measurements for all but the $z \\sim 0.1$ galaxies are based on HST imaging, and SFRs are derived using a Herschel-calibrated ladder of SFR indicators.  We find that a correlation between the structure and stellar population of galaxies (i.e., a 'Hubble sequence') is already in place since at least $z \\sim 2.5$.  At all epochs, typical star-forming galaxies on the main sequence are well approximated by exponential disks, while the profiles of quiescent galaxies are better described by de Vaucouleurs profiles.  In the upper envelope of the main sequence, the relation between the SFR and Sersic index reverses, suggesting a rapid build-up of the central mass concentration in these starbursting outliers.  We observe quiescent, moderately and highly star-forming systems to co-exist over an order of magnitude or more in stellar mass.  At each mass and redshift, galaxies on the main sequence have the largest size.  The rate of size growth correlates with specific SFR, and so does $\\Sigma_{SFR}$ at each redshift.  A simple model using an empirically determined SF law and metallicity scaling, in combination with an assumed geometry for dust and stars is able to relate the observed $\\Sigma_{SFR}$ and $SFR_{IR}/SFR_{UV}$, provided a more patchy dust geometry is assumed for high-redshift galaxies. ",
        "introduction": "\\label{obs_method_sample.sec}  In our analysis, we divide the galaxies in three redshift bins: $0.02 < z < 0.2$, $0.5 < z < 1.5$, and $1.5 <  z < 2.5$, which we will in short refer to as $z \\sim 0.1$, $z \\sim 1$, and $z \\sim 2$.  The $z \\sim 0.1$ sample is extracted from the Sloan Digital Sky Survey (SDSS), as detailed in Section\\ \\ref{SDSS_GALEX.sec}.  The higher redshift samples are extracted from four deep fields that host some of the richest multi-wavelength data sets available to date: COSMOS, UDS, GOODS-South, and GOODS-North.  We describe the data available for them in Sections\\ \\ref{COSMOS.sec} to \\ref{GOODSN.sec}, and summarize the methods applied to compute derived properties such as photometric redshifts, stellar masses, SFRs and structural parameters in Section\\ \\ref{derived_prop.sec}.  \\subsection{Fields and Data} \\label{fields_data.sec}  \\subsubsection{SDSS + GALEX} \\label{SDSS_GALEX.sec}  SDSS offers a robust local universe anchor for our study of galaxy properties in the SFR-Mass diagram, and its evolution with redshift.  We use the spectroscopic SDSS sample (DR7, Abazajian et al. 2009) in overlap between the MPA-JHU and NY-VAGC value-added galaxy catalogs.  SFRs and stellar masses were taken from the MPA-JHU compilation\\footnote{http://www.mpa-garching.mpg.de/SDSS/DR7}.  Briefly, the stellar masses were estimated  by fitting Bruzual \\& Charlot (2003, hereafter BC03) stellar population synthesis models with a wide range of star formation histories (SFHs) to the broad-band photometry (Salim et al. 2007).  Unless explicitly stated otherwise, we use total, dust- and aperture-corrected SFRs.  SFRs within the SDSS spectroscopic fiber were primarily based on H$\\alpha$, using the calibration of Brinchmann et al. (2004), which includes an extinction correction based on the Balmer decrement.  Aperture corrections to total use fits to the photometry of the outer regions of the galaxies, as described by Salim et al. (2007).  We use the SDSS (DR7) - GALEX (GR5) matched catalog from Bianchi et al. (2011) to compute the unobscured part of the star formation $SFR_{UV}$.  Sizes and profile shapes were measured by the NY-VAGC team\\footnote{http://sdss.physics.nyu.edu/vagc} by means of Sersic fits to the galaxy surface brightness profiles (Blanton et al. 2003, 2005).  By default, we adopt the structural parameters measured in the $g$ band, but we will discuss the dependence on wavelength where appropriate.  \\subsubsection{COSMOS} \\label{COSMOS.sec}  We make use of public multi-wavelength photometry from Ilbert et al. (2009) and Gabasch et al. (2008) over an effective area of 1.48 deg$^2$ in the COSMOS HST field (Scoville et al. 2007; Koekemoer et al. 2007).  Our working catalog is cut at $i < 25$ to guarantee sufficient signal-to-noise to allow for a reliable determination of photometric redshifts and other derived properties.  A total of 36 medium and broad bands sample the SED from GALEX to IRAC wavelengths.  At infrared wavelengths, PACS imaging from PEP, extracted using 24 $\\mu$m sources as prior, reaches depths of 5.7 and 11.9 mJy (3$\\sigma$) at 100 $\\mu$m and 160 $\\mu$m respectively.  MIPS 24 $\\mu$m imaging in COSMOS (Sanders et al. 2007; Le Floc'h et al. 2009) is used down to 60 $\\mu$Jy.  An X-ray catalog based on XMM-Newton observations is released by Cappelluti et al. (2009).  High-resolution $I_{814}$-band imaging with HST/ACS ($I_{814} < 27,\\ 5\\sigma$) spanning the entire field, and probing the rest-frame $B$ and $2700\\AA$ band at $z \\sim 1$ and $z \\sim 2$ respectively, forms the foundation for structural measurements in COSMOS.  \\subsubsection{UDS} \\label{UDS.sec}  A deep WFC3 $H_{160}$-selected catalog ($H_{160} < 26.7$, 5$\\sigma$) with consistent photometry in 16 bands from $B$ to 8 $\\mu$m was constructed on the basis of CANDELS, SEDS (PI G. Fazio), SpUDS (PI J. Dunlop), and ancillary ground-based data in the $\\sim 200$ arcmin$^2$ CANDELS-covered UDS area (Guo et al. in prep).  MIPS 24 $\\mu$m imaging to a depth of 30 $\\mu$Jy (3$\\sigma$) was carried out as part of the SpUDS survey.  Ueda et al. (2008) presented the X-ray source catalog of the Subaru/XMM-Newton deep survey, in which the CANDELS/UDS area is embedded.  We performed the structural measurements on the WFC3 $H_{160}$ mosaic, which with 4/3 orbits per pointing forms part of CANDELS-Wide.  The $H_{160}$ band probes the rest-frame $I$ and $V$ band at $z \\sim 1$ and $z \\sim 2$ respectively.  \\subsubsection{GOODS-South} \\label{GOODSS.sec}  Following similar procedures as for UDS, Guo et al. (in prep) constructed a deep WFC3 $H_{160}$-selected catalog ($H_{160} < 27$, 5$\\sigma$) for the Early Release Science (ERS, PI O'Connell) and CANDELS-Deep area in GOODS-South.  The catalog contains consistent photometry in 14 passbands from $U$ to 8 $\\mu$m.  Spanning the same wavelength range, we use FIREWORKS ($K_s < 24.3,\\ 5\\sigma$, Wuyts et al. 2008) as supporting multi-wavelength catalog to exploit CANDELS-Wide data in the bottom 20\\% of the 148 sq. arcmin GOODS-South field.  As part of the PEP survey (Lutz et al. 2011), PACS photometry was obtained to a 5$\\sigma$ depth of 1.8 mJy, 1.9 mJy and 3.3 mJy at 70 $\\mu$m, 100 $\\mu$m and 160 $\\mu$m respectively, using the position of 24 $\\mu$m sources as prior.  The 24 $\\mu$m imaging itself (30 $\\mu$Jy, 5$\\sigma$) was taken from Magnelli et al. (2009).  Where appropriate, we select X-ray sources from the Chandra 2 Ms source catalog by Luo et al. (2008).  We performed the structural measurements on WFC3 $H_{160}$-band imaging drizzled to a $0\\farcs06 / pixel$ scale, obtained as part of the ERS and CANDELS programs.  The WFC3/IR observations of the ERS were obtained in F098M, F125W and F160W for a total depth of 2 orbits per pointing in each filter (Windhorst et al. 2011) and the mosaics used here were produced using MultiDrizzle (Koekemoer et al. 2002) following the approach outlined in Koekemoer et al. (2011).  The 4-epoch CANDELS-Deep area is similar in depth, while CANDELS-Wide received half the integration time.  \\subsubsection{GOODS-North} \\label{GOODSN.sec}  We derived photometric redshifts and stellar masses for galaxies in GOODS-North based on a $z+K$-selected catalog with photometry in 16 bands from GALEX to IRAC wavelengths (Berta et al. 2010).  Additional longer wavelength information at 24 $\\mu$m (30 $\\mu$Jy, 5$\\sigma$), 100 $\\mu$m (5.1 mJy, 5$\\sigma$), and 160 $\\mu$m (8.7 mJy, 5$\\sigma$) was acquired as part of the GOODS and PEP surveys.  Where appropriate, we select X-ray sources from the Chandra 2 Ms source catalog by Alexander et al. (2003).  Structural measurements in GOODS-North are based on the GOODS ACS $z_{850}$-band imaging (Giavalisco et al. 2004), corresponding to the rest-frame $B$ and $3000\\AA$ band at $z \\sim 1$ and $z \\sim 2$ respectively..  \\subsection{Derived Galaxy Properties} \\label{derived_prop.sec}  For all deep fields, we use identical procedures to derive photometric redshifts, stellar masses, SFRs, and structural parameters, hence optimizing the consistency of the combined data set.  \\subsubsection{Photometric Redshifts}  We determined photometric redshifts by fitting a superposition of 6 galaxy templates to the observed broad-band SEDs using the photometric redshift code EAZY (Brammer et al. 2008).  For the 6955 out of 132328 (5.3\\%) of galaxies at $0.5 < z < 1.5$, and 497 out of 35649 (1.4\\%) of galaxies at $1.5 < z < 2.5$ that have a spectroscopic redshift from one of the many spectroscopic campaigns in the above four fields (see, e.g., Vanzella et al. 2008; Lilly et al. 2009; Barger et al. 2008), we adopt the $z_{spec}$ as redshift in our analysis.  Comparing the photometric redshifts $z_{phot}$ and spectroscopic redshifts $z_{spec}$ of spectroscopically confirmed galaxies, we find a median and scatter (normalized median absolute deviation) in $\\Delta z / (1+z)$ of (-0.001; 0.015) in the $z \\sim 1$ and (-0.007; 0.052) in the $z \\sim 2$ redshift bin.  Naturally, these uncertainties only apply to galaxies with similarly bright magnitudes as represented in the spectroscopic sample based on which the scatter was computed.  E.g., the public zCOSMOS-bright spectroscopic survey (Lilly et al. 2009) reaches down to $i = 22.5$.  For this spectroscopic sample in COSMOS, we measure a scatter in $\\Delta z / (1+z)$ of 0.010.  As detailed in Appendix B.3, the estimated uncertainty increases to $\\sigma_{NMAD} (\\Delta z / (1+z)) = 0.048$ for COSMOS galaxies with $24 < i < 25$, which account for 40\\% of the galaxies in that field, and 30\\% of our entire sample.  The estimated uncertainty is consistent with earlier work by Ilbert et al. (2009).  For the other deep fields as well, the increase in the photometric redshift uncertainty with respect to the spectroscopic sample amounts to a factor of $\\sim 5$ for the faintest galaxies that enter our analysis.  An extensive analysis in Appendix B.3 shows that our results are robust against photometric redshift uncertainties, even when accounting for the fact that the $z_{phot}$ quality of the entire sample is likely lower than that of the (typically brighter) spectroscopically confirmed subsample.  \\subsubsection{Stellar Masses} \\label{SEDmodeling.sec}  We used FAST (Kriek et al. 2009a) to fit BC03 models with exponentially declining SFHs to the $\\lambda_{obs} \\leq 8\\ \\mu$m broad-band SEDs of the galaxies.  Following the recipe presented by Wuyts et al. (2011), we impose a minimum e-folding time of 300 Myr.  The time since the onset of star formation was allowed to vary between 50 Myr and the age of the universe at the observed redshift.  Our choice of SFHs guarantees consistency with most of the literature, and in addition has the advantage (as opposed to, e.g., only exponentially increasing SFHs) that it is also applicable to lower redshift galaxies and high-redshift quiescent galaxies, that clearly formed stars at a higher rate in their past than at the epoch of observation.  Maraston et al. (2010) showed that adopting exponentially increasing SFHs for star-forming galaxies at $z \\sim 2$ leads to negligible changes in the estimated stellar masses, unless also a stringent constraint on formation redshift is imposed, in which case the estimated stellar mass of these galaxies can increase by a few 0.1 dex.  Solar metallicities were assumed throughout, and dust attenuation was modeled using the Calzetti et al. (2000) reddening law, with visual extinctions in the range $0 < A_V < 4$.  For a detailed description of the dependence of estimated stellar masses on the assumed metallicity and attenuation law, we refer the reader to Wuyts et al. (2007).  We revisit the assumption of solar metallicity in Section\\ \\ref{model_IRUV.sec} in the context of a model for IR/UV ratios that involves a metallicity scaling, and conclude that it has a negligible impact on our results.  As for the SDSS galaxies, we assumed a universal Chabrier (2003) IMF for all galaxies in the lookback surveys.  \\subsubsection{Star Formation Rates}  Exploiting deep PACS photometry in the GOODS-South field, Wuyts et al. (2011) established a continuity across SFR indicators, ranging from UV + FIR, over UV + MIR, to SED modeled SFRs.  We adopt this cross-calibrated 'ladder of SFR indicators' to obtain an estimate of the SFR for each galaxy in our sample individually.  Briefly, if the galaxy is detected in an IR band, we compute the total SFR by summing the unobscured and re-emitted emission from young stars, following Kennicutt (1998):  \\begin {equation} SFR_{UV + IR}\\ [M_{\\sun} yr^{-1}] = 1.09 \\times 10^{-10}\\ (L_{IR} + 3.3 L_{2800})/L_{\\sun} \\label{SFRuvir.eq} \\end {equation}  where $L_{2800} \\equiv \\nu L_{\\nu}(2800\\AA)$ was computed with EAZY from the best-fitting SED.  Here, the total IR luminosity $L_{IR} \\equiv L(8-1000\\mu m)$ is derived monochromatically from the longest wavelength IR band that is significantly ($> 3\\sigma$) detected: either PACS 160, 100, or 70 $\\mu$m, or MIPS 24 $\\mu$m.  We use the mid- to far-infrared SED template by Wuyts et al. (2008) for the conversion to $L_{IR}$.  This conversion leads to a consistency between 24 $\\mu$m and PACS-derived $L_{IR}$, unlike locally calibrated template sets that have been proven to overestimate the $L_{IR}$ of $z \\gtrsim 2$ galaxies based on their rest-frame 8 $\\mu$m emission (Nordon et al. 2010; Elbaz et al. 2010; Wuyts et al. 2011).  For galaxies that lack an IR detection, we adopt the SFR associated to the best-fit stellar population synthesis model from Section\\ \\ref{SEDmodeling.sec}.  For the low- to intermediate SFR regime, where $SFR_{UV+IR}$ is only attainable in the deepest of our four fields (GOODS-South), Wuyts et al. (2011) demonstrated that SED modeled SFRs are consistent with those obtained from Equation\\ \\ref{SFRuvir.eq}.  As for the stellar masses, the values of $SFR_{SED}$, obtained without imposing explicit constraints on the formation redshift, would not change significantly when assuming exponentially increasing rather than decreasing SFHs (Maraston et al. 2010).  \\subsubsection{Structural Parameters} \\label{structural_params.sec}  In this paper, we limit our analysis of galaxy structure to a parametric approach, focussing on two basic properties of the surface brightness distribution: its extent and cuspiness.  The former is expressed as the circularized effective radius $r_e \\equiv r_{e, major} \\times \\sqrt{b/a}$ corresponding to the best-fit Sersic profile.  The latter is quantified by the Sersic index $n$, with $n$ restricted to the range $0.2 < n < 8$.  We adopt the structural measurements performed on the longest wavelength high-resolution imaging available for each field (WFC3 for GOODS-S and UDS, ACS for GOODS-N and COSMOS, see Section\\ \\ref{fields_data.sec}).  The GALAPAGOS package (Barden et al. 2009) or equivalent procedures were used to pre-determine the sky background and the neighboring sources that need to be simultaneously fit or masked, and to automate fitting Sersic profiles to the 2D surface brightness distributions with GALFIT (C. Y. Peng et al. 2010).  GALFIT takes into account convolution by the point spread function (PSF).  \\begin {figure*}[t] \\centering \\plotone{f1.eps} % \\caption{ Surface brightness profile shape in the SFR-Mass diagram.  A 'structural main sequence' is clearly present at all observed epochs, and well approximated by a constant slope of 1 and a zeropoint that increases with lookback time ({\\it white line}).  While SFGs on the MS are well characterized by exponential disks, quiescent galaxies at all epochs are better described by De Vaucouleurs profiles.  Those galaxies that occupy the tip and upper envelope of the MS also have cuspier light profiles, intermediate between MS galaxies and red \\& dead systems. \\label{n.fig}} \\vspace{-0.12in} \\end {figure*}  Since the longest wavelength at which high-resolution imaging is presently available varies between our four deep fields, our study lacks the potential to consistently probe the rest-frame optical regime, and fully control morphological k-corrections, if present.  However, since we are interested not only in galaxies on the main sequence, but also in those rare systems that form the high-SFR tail, we prefer to sacrifice this aspect of our analysis, rather than area.  As a sanity check, we compared sizes and Sersic indices measured in the $z_{850}$ and $H_{160}$ imaging of GOODS-S (see Appendix B.2).  We find the median ACS - WFC3 deviations to be small (-0.01 in $\\Delta \\log r_e$ and -0.10 in $n$ at $1.5 < z < 2.5$) relative to the trends that will be discussed in this paper, suggesting that any biases from morphological k-corrections have only a limited impact on our conclusions.  Bond et al. (2011) arrive at a similar conclusion comparing rest-frame optical and rest-frame UV size measurements of $z \\sim 2$ galaxies.  Moreover, we also verified that the results obtained for each of the four fields individually are consistent with those presented for the combined data set in this paper, albeit with more noise due to the smaller number statistics (see Appendix B.1).  While beyond the scope of this paper, the wavelength dependence of structural properties can potentially reveal interesting clues on the physical processes shaping galaxies in the young universe (see, e.g., Guo et al. 2011; Szomoru et al. 2011; and Cassata et al. 2011, who find evidence for mild morphological k-corrections such that the centers of galaxies in their samples tend to be somewhat redder).  For an in depth discussion of resolved stellar populations inferred from spatial variations in color, we defer the reader to Wuyts et al. (in prep).  \\subsection{Sample Selection} \\label{sample.sec}  Our four deep fields are not uniform in depth.  Consequently, they have different completeness limits in the SFR-Mass diagram.  Since we are mainly interested in individual galaxy characteristics rather than abundances (number or mass densities) as function of position along or across the MS, we refrain from applying any incompleteness corrections to the observed populations.  Under the premise that, at any given redshift, galaxies form a two-parameter family described by their mass and SFR, this approach works well.  I.e., a median galaxy property can be computed reliably, and unbiased by any completeness issues, based on the objects observed in a given bin of SFR-Mass space.  Our final sample comprises 639924 galaxies at $0.02 < z < 0.2$, 132328 galaxies at $0.5 < z < 1.5$, and 35649 galaxies at $1.5 < z < 2.5$.  The relative breakdown in galaxies of different masses is determined by the depth of the observations, and the stellar mass function at the respective redshifts.  Above $M > 10^{10}\\ M_{\\sun}$ our sample counts 532131, 31127, and 8895 galaxies at $z \\sim 0.1$, $z \\sim 1$, and $z \\sim 2$ respectively.  Above $M > 10^{11}\\ M_{\\sun}$, the numbers drop to 147922, 2767, and 1059 galaxies at $z \\sim 0.1$, $z \\sim 1$, and $z \\sim 2$ respectively.  An overview of the sample size per field is provided in Table\\ \\ref{fields.tab}.  \\section {Results on Galaxy Structure} \\label{structure.sec}  \\subsection {Profile Shape} \\label{profile.sec}  \\begin {figure*}[t] \\centering \\plotone{f2.eps} % \\caption{ {\\it Top panels:} SFR histograms relative to the 'structural main sequence' for galaxies in the mass range $5 \\times 10^{9}\\ M_{\\sun} < M < 10^{11}\\ M_{\\sun}$.  {\\it Bottom panels:} Sersic index as function of deviation from the MS, measured along the SFR axis, for the same $10^{9.7} - 10^{11}\\ M_{\\sun}$ mass slice.  The black curve and error bars indicate the median in bins of $\\Delta \\log (SFR)$, and respective errors in the median.  The grey polygon marks the central 68th percentile of the distribution.  Colored dotted curves in the left-hand panel illustrate how the $z \\sim 0.1$ relation changes when adopting structural measurements performed at shorter/longer wavelengths than the SDSS $g$ band.  Relative to the MS, a correlation between profile shape and star-forming activity is present and looks similar at all epochs.  The relation is not monotonic, but rather shows a reversal at the high-SFR end. \\label{dev.fig}} \\vspace{-0.12in} \\end {figure*}  We start by analyzing the surface brightness profile shape as a function of position in the SFR-Mass diagram in Figure\\ \\ref{n.fig}.  The three panels show from left to right the $z \\sim 0.1$, $z \\sim 1$, and $z \\sim 2$ bins respectively.  Instead of indicating the relative abundance of galaxies in different regions of the diagram, we use the color-coding to mark the median value of the Sersic index $n$ of all galaxies in each $[SFR, M]$ bin.  For displaying purposes, we restrict the range of the colorbar to $1<n<4$, and assign the same color as $n=1$ and $n=4$ to bins with median $n < 1$ or median $n > 4$ respectively.  The fractions $(f_{n<1};\\ f_{n>4})$ of galaxies lying outside these bounds amounts to (0.09; 0.24), (0.41; 0.14), and (0.41; 0.16) at $z \\sim 0.1$, $z \\sim 1$, and $z \\sim 2$ respectively.  The fraction of $[SFR, M]$ bins with median n outside this range is small: (0.02; 0.11) at $z \\sim 0.1$, (0.14; 0.15) at $z \\sim 1$, and (0.11; 0.11) at $z \\sim 2$.  The resulting diagrams present a remarkably smooth variation in the typical galaxy profile shape across the diagram.  Moreover, despite the loss of information on number densities, the so-called main sequence of star formation is immediately apparent, and its presence persists out to the highest observed redshifts.  This 'structural main sequence' consists of galaxies with near-exponential profiles ($n \\approx 1$) and shows a similar behavior as the conventional 'number main sequence' as identified on the basis of number densities in the SFR-Mass diagram (e.g., Noeske et al. 2007; Elbaz et al. 2007; Daddi et al. 2007).  Namely, an upward shift of the zeropoint is observed with increasing lookback time.  At each epoch, the MS in Figure\\ \\ref{n.fig} is well approximated by a slope of unity ({\\it white line}).  The SFR at which the median n reaches a minimum in a mass slice around $\\log(M) = 10$ roughly coincides with the mode of the $\\log(SFR)$ distribution in that mass slice, but depending on the fitting method and sample definition used to weed out quiescent galaxies, a somewhat shallower slope than unity may be measured for the 'number main sequence' at the massive end (see, e.g., Rodighiero et al. 2010).  Below the structural MS, a cloud of galaxies with cuspy, near-De Vaucouleurs ($n \\approx 4$) profiles is visible.  This population of massive quiescent galaxies is present at all observed epochs.  Our first and foremost conclusion from Figure\\ \\ref{n.fig} is therefore that a correlation between the structure and stellar populations of galaxies is already in place since $z \\sim 2.5$.  In the local universe, the presence of such a correlation has been described in depth by, e.g., Kauffmann et al. (2003) and Brinchmann et al. (2004).  Scarlata et al. (2007) demonstrated that photometrically selected massive quiescent galaxies are already dynamically relaxed (i.e., have the morphological appearance of early-type galaxies) at $z \\sim 1$.  Also out to $z \\sim 1$, Maier et al. (2009) exploited the spectroscopic zCOSMOS survey to reveal a correlation between the specific SFR of galaxies and the Sersic index of their surface brightness profiles.  At larger lookback times, our results are in agreement with previous findings, based on smaller samples, reporting a similarly early emergence of the Hubble sequence on the basis of stellar surface mass density, galaxy size and visual appearance (Franx et al. 2008; Toft et al. 2009; Kriek et al. 2009b).  We find that, across cosmic time, the typical Sersic index of galaxies is not optimally described as a function of their absolute SFR, or even their absolute specific SFR, but rather as a function of their position relative to the MS at the epoch of observation.  The correspondence between mass, SFR and structure, as quantified by the Sersic index, is equivalent to the Hubble sequence.  Based on samples of unprecedented size, we see that such a sequence already existed at $z \\sim 2$; bulge-dominated morphologies go hand in hand with a more quiescent nature.  First indications of such a correlation between Sersic index and specific SFR out to high redshift were recently reported by Szomoru et al. (2011) based on 27 galaxies at $z \\sim 1$ and 16 galaxies at $z \\sim 2$.  \\begin {figure*}[t] \\centering \\plotone{f3.eps} % \\caption{ Galaxy size in the SFR-Mass diagram.  Moving up the main sequence, towards the high-mass end, galaxies grow bigger (i.e., a size-mass relation is observed at all epochs).  With respect to both quiescent galaxies, and also the high-SFR tail, galaxies on the MS have the largest size.  Finally, in any given $[\\log (M); \\log (SFR)]$ bin, galaxies grow with time.  Quantitatively, this growth is strongest for quiescent (low $SFR/M$) galaxies. \\label{re.fig}} \\vspace{-0.12in} \\end {figure*}  In detail, an additional interesting aspect is revealed by the $n$ distribution in SFR-Mass space.  The surface brightness profile shape does not vary monotonically across the main sequence, but instead shows a reversal in the high-SFR tail of the distribution, such that the upper envelope of the MS is composed of galaxies whose typical cuspiness is intermediate between those of normal MS galaxies and the quiescent population below the MS.  This is illustrated further in Figure\\ \\ref{dev.fig}, where we plot a cut through of the $5 \\times 10^{9}\\ M_{\\sun} < M < 10^{11}\\ M_{\\sun}$ mass slice.   The top panels show the histogram of deviations from the MS measured along the SFR axis: $\\Delta \\log(SFR)$.  The top left panel illustrates the well-known bimodality between star-forming and passive galaxies in the local universe (e.g., Kauffmann et al. 2003).  A similar bimodality, albeit with a reduced amplitude of the quiescent peak, is visible at $z \\sim 1$ (and becomes more pronounced if we were to limit ourselves to higher masses only).  While both populations are also found in our highest redshift bin, the quiescent fraction is further reduced.  Incompleteness in our shallowest field, as well as uncertainties in photometric redshifts and estimated SFRs inhibit the identification of a true bimodality in the $z \\sim 2$ histogram.  The bottom panels of Figure\\ \\ref{dev.fig} show the median dependence of $n$ on $\\Delta \\log(SFR)$ ({\\it black curve}), as well as errors on the determination of the median ({\\it black error bars}).  With grey polygons, we mark the central 68th percentile.  Clearly, a variety of profile shapes is measured, even at a given $\\Delta \\log (SFR)$, and particularly in the $z \\sim 1$ and $z \\sim 2$ bins.  This stems partly from measurement uncertainties, that dominate in the low surface brightness regime, but in addition likely means that in detail galaxies are more complex than a simple two-parameter ($[SFR, M]$) family.  Interestingly, the scatter seems to be minimized for galaxies that reside on the MS.  F\\\"{o}rster Schreiber et al. (2011) present a detailed analysis of large and clumpy MS galaxies at $z \\sim 2$, that exhibit shallow surface brightness profiles ($n < 1$).  Such a population is also present in our sample (see Figure\\ \\ref{dev.fig} and the above quoted $f_{n<1}$).  Focussing on the overall trend, we find the results at $z \\sim 1$ and $z \\sim 2$ to be very similar to those established at $z \\sim 0.1$.  Namely, galaxies on the MS resemble exponential disks, high-SFR outliers have a centrally enhanced surface brightness profile, and once galaxies end up below the MS, they quickly reach a plateau of $n \\approx 4$.  We discuss the physical implications of this empirical relation in Section\\ \\ref{quenching.sec}.  We verified that, when excluding X-ray sources from our analysis, the same median dependence of the surface brightness profile shape on the location of galaxies in the SFR-Mass diagram remains present.  This is the case even for the fields with the deepest X-ray data.  The same robustness also applies to the other galaxy properties discussed later in this paper.  Our test indicates that the observed trends (particularly the enhanced Sersic indices of the high-SFR outliers at a given mass) is not merely due to a nuclear, non-stellar point source.  When focussing on the X-ray detected sources alone, we find that they occupy the mass regime above $\\log M \\gtrsim 10$, and span a wide range in SFRs, from the quiescent class to on and above the MS.  Within a given bin, their surface brightness profiles differ in the sense that they have higher Sersic indices and smaller sizes than normal galaxies of the same $[SFR, M]$.  Point source contamination would drive the structural parameters in the observed direction, but extensive simulations (see, e.g., Simmons \\& Urry 2008) are required to constrain whether in addition there is any evidence for a different stellar structure of the hosts.  Before moving to the next structural property, we have a closer look at the $z \\sim 0.1$ bin in Figure\\ \\ref{dev.fig}, where we use structural measurements in five passbands to investigate how the observed trend depends on wavelength.  The dotted colored curves in the bottom left panel of Figure\\ \\ref{dev.fig} illustrate the median relation derived from Sersic fits in the SDSS $u$, $r$, $i$, and $z$ bands.  Together with the default $g$-band profile fits ({\\it solid black curve}), a trend emerges in which the overall relation between $n$ and $\\Delta \\log (SFR)$ is remarkably robust against morphological k-corrections.   The largest difference is observed for galaxies that lie on the MS ($\\Delta \\log (SFR) = 0$), such that a cuspier profile is inferred from redder waveband imaging.  This comes as no surprise, as it is well known that many normal SFGs in the nearby universe are composed of a (red, $n \\approx 4$) bulge and (blue, $n \\approx 1$) disk component\\footnote{The relative fraction of such bulged disks is observed to decrease with redshift, at least out to $z \\sim 1$ (Sargent et al. 2007).}.  At low and high values of $\\Delta \\log (SFR)$, the Sersic index levels of at similar values ($n = 4$ and $n = 2$ respectively), irrespective of wavelength.  This suggests that the light concentration in the centers of these galaxies corresponds to a mass concentration, and is not due (only) to the large luminosity of a young stellar population.  \\subsection {Size} \\label{size.sec}  \\begin {figure*}[t] \\centering \\plotone{f4.eps} % \\caption{ Surface density of star formation ($\\Sigma_{SFR}$) in the SFR-Mass diagram.  Lines of equal $\\Sigma_{SFR}$ run approximately parallel to the MS, as do lines of constant $SFR/M$. \\label{surf.fig}} \\vspace{-0.12in} \\end {figure*}  Having established the existence of a 'Hubble sequence' out to $z \\sim 2.5$, we now turn to the zeroth order structural measurement of size, and address how it varies along and across the MS at each of our observed epochs.  Figure\\ \\ref{re.fig} is similar in nature to Figure\\ \\ref{n.fig}, but the color-coding is now used to mark differences in the spatial extent of the surface brightness distribution ($r_e$).  As in Figure\\ \\ref{n.fig}, a qualitatively similar behavior is observed at all redshifts.  Overall, more massive galaxies tend to be larger.  However, the mass-size relation is not fundamental in the sense that at a given mass substantial variations in size occur that are not random, but instead correlate with the offset from the MS (see also the Appendix, Figure\\ \\ref{dev_app.fig}).  This translates to different mass-size relations followed by passive and star-forming galaxies, as reported previously by Shen et al. (2003) based on SDSS, but also by Williams et al. (2010) and Szomoru et al. (2011) who extended the analysis for massive galaxies out to $z = 2$, based on ground-based $K$-band imaging in UDS and WFC3 $H_{160}$-band imaging in the Hubble Ultra Deep Field respectively.  Consistent with previous results by, e.g., Franx et al. (2008) and Toft et al. (2009), high-redshift quiescent galaxies are more compact than the bulk of SFGs at the same epoch.  Also at $z \\sim 0.1$, galaxies below the MS (at least at $M < 10^{11}\\ M_{\\sun}$) tend to be smaller than those on the MS.  The high-SFR tail of the $z \\sim 0.1$ distribution differs from normal MS galaxies in the sense that they tend to have higher surface mass densities (i.e., are more compact for a given mass).  As we march up in redshift, our measurements are consistent with a similar behavior across all lookback times considered.  Comparing the panels for the three redshift bins, it is immediately apparent that galaxies grow over time.  We quantified the growth by fitting a $r_e(z) / r_e(z=0) = (1+z)^\\alpha$ relation to the galaxies in each $[SFR, M]$ bin.  It is important to note that this reflects the size evolution of the galaxy population as a whole, and not per se that of any given galaxy individually, as galaxies will move from one $[SFR, M]$ bin to another as they build up stellar mass, and new systems will be added over time.  We find $\\alpha$ to be varying across the SFR-Mass diagram, from $\\alpha = -1.2_{-0.4}^{+0.9}$ below $\\log (SFR/M) < -11$ to $\\alpha = -0.4_{-0.7}^{+0.7}$ above $\\log (SFR/M) > -11$.   Here, the range indicates the normalized median deviation over the different $[SFR, M]$ bins.  Splitting the latter class into galaxies above and below the MS at $z \\sim 0.1$ ({\\it white line in the left-hand panel of Figure\\ \\ref{re.fig}}), we obtain $\\alpha = -0.2_{-0.6}^{+0.5}$ and $\\alpha = -0.9_{-0.5}^{+0.6}$ respectively.  In other words, the rate of size evolution correlates with the specific star formation rate, reaching the lowest values in the (especially massive) quiescent population. ",
        "conclusions": ""
    },
    "1107/1107.2548_arXiv.txt": {
        "abstract": "Helicity and $\\alpha$ effect  driven by the nonaxisymmetric  Tayler instability of toroidal magnetic fields in  stellar radiation zones are computed. In the linear approximation a purely toroidal field always excites  pairs of modes  with identical   growth rates but with opposite helicity so that the net helicity  vanishes. If the magnetic background field has a helical structure by an extra  (weak) poloidal component then one of the modes dominates producing a net kinetic helicity anticorrelated to the current helicity of the background field.\\\\ The mean electromotive force is computed with the  result that the  $\\alpha$ effect  by   the most rapidly growing  mode has the same sign as the current helicity of the background field. The  $\\alpha$ effect is found as  too small to drive an $\\alpha^2$ dynamo but  the excitation conditions for an $\\alpha\\Omega$ dynamo can     be fulfilled  for  weak poloidal fields. Moreover, if  the dynamo produces its own $\\alpha$ effect by the magnetic instability then   problems with its  sign do not arise. For all cases,  however,  the $\\alpha$ effect    shows an extremely  strong concentration to the poles so that a possible $\\alpha\\Omega$ dynamo might only work at the polar regions. Hence,   the  results of our linear theory lead to a new topological problem for the existence of  large-scale dynamos in  stellar radiation zones on the basis of the current-driven instability of toroidal fields. ",
        "introduction": "Hydromagnetic dynamos can be understood   as magnetic instabilities driven by  a special flow pattern in fluid conductors.  There are, however, strong restrictions on the characteristics  of such flows (see Dudley \\& James 1989) as well as on the geometry of the resulting magnetic fields  \\citep{C33}. The restrictions even exclude any dynamo activity  for a number of flows.  We mention as  an example that differential rotation alone can never maintain a dynamo (Elsasser 1946).  An open question is whether  magnetic instabilities are able to  excite   a sufficiently complicated motion that together with a (given)  background flow can generate magnetic fields. \\citet{TP92} suggested that nonuniformly rotating disks can produce a dynamo when magnetorotational (MRI) and magnetic buoyancy instabilities are active. Later on, numerical simulations of \\citet{Bea95} and \\citet{Hea96} have shown that MRI alone may be sufficient for the accretion disk dynamo. It remains, however, to check (at least for the case of low magnetic Prandtl number) whether the MRI-dynamo has physical or numerical origin \\citep{FP07,Fea07}.  Another possibility was discussed by \\citet{S02} who suggested that differential rotation and magnetic kink-type instability \\citep{T73} can jointly drive a dynamo in stellar radiation zones. The dynamo if real would be very important for the angular momentum transport in stars and their secular evolution. It taps energy from differential rotation thus reducing the rotational shear. Radial displacements converting toroidal magnetic field into poloidal field are necessary for the dynamo. The dynamo, therefore, unavoidably mixes chemical species in stellar interiors that may have observable consequences for stellar evolution.  Such a dynamo, however,  has not yet been  demonstrated to exist. The doubts especially concern  the kink-type instability that in contrast to MRI exists also without differential rotation. The Tayler instability develops in expense of magnetic energy. Estimations of dynamo parameters are thus necessary to assess  the dynamo-effectiveness of this magnetic instability.  The basic role in turbulent dynamos plays  the ability of correlated magnetic (${\\vec b}$) fluctuations and velocity (${\\vec u}$) fluctuations to produce a mean electromotive force along the background magnetic field ${\\vec B}$ and also along the electric current $\\vec J$, i.e. \\begin{equation} \\langle {\\vec u}\\times{\\vec b}\\rangle = \\alpha{\\vec B}- \\mu_0 \\eta_{\\rm T}{\\vec J} . \\label{1} \\end{equation} We estimate   the $\\alpha$ effect by Tayler instability in the present paper. We do also find   indications for the appearance of the turbulent diffusivity $\\eta_{\\rm T}$ in the calculations but  we do not follow them here in detail. For purely toroidal fields we did {\\em not} find indication for the existence of the term ${\\vec \\Omega}\\times \\vec{J}$ which can appear in the expression (\\ref{1}) in form  of a rotationally induced anisotropy of the diffusivity tensor.   The fluctuating fields for the most rapidly growing eigenmodes and the azimuthal averaging are applied in the LHS of Eq.\\,(\\ref{1}) to estimate the $\\alpha$ effect and its relation to the kinetic and magnetic helicity ${\\cal H}^{\\rm kin}$ and ${\\cal H}^{\\rm curr}$. Our linear stability computations do not allow the evaluation of the $\\alpha$ effect amplitude but its latitudinal profile  and its ratio to the product of rms values of $\\vec{u}$ and $\\vec{b}$ (i.e. the  correlation coefficient) can be found. As the differential rotation is necessary for dynamo, we estimate also the influence of differential rotation  on Tayler instability. Next, a dynamo model with the parameters estimated for the magnetic instability is designed to find the global modes of the instability-driven dynamo. ",
        "conclusions": "It is demonstrated with a linear theory that the Tayler instability of a toroidal magnetic field in a density-stratified radiation zone of hot stars does not produce helicity and/or $\\alpha$ effect. If, however, a weak poloidal field component is added forming a large-scale current helicity ${\\vec B} \\cdot {\\rm curl} {\\vec B}$ then the instability leads to small-scale helicity, current helicity and also to $\\alpha$ effect. If ${\\vec B} \\cdot {\\rm curl} {\\vec B}$ is {\\em positive} in the northern hemisphere then the helicity $\\langle {\\vec u}\\cdot {\\rm curl} {\\vec u} \\rangle$ is {\\em negative} at the northern hemisphere if only  the modes are considered which grow fastest. Hence, the small-scale kinetic helicity and the large-scale current helicity of the background field are anticorrelated.  As it is often the case, the corresponding $\\alpha$ effect is anticorrelated with the kinetic helicity and, therefore, positively correlated with the large-scale current helicity ${\\vec B} \\cdot {\\rm curl} {\\vec B}$. The calculations lead to two basic properties of this $\\alpha$ effect. It is i) small, i.e. the normalized value $C_\\alpha$ is only of order $10^{-(2\\dots 3)}$ and ii)  concentrated at the poles. The first property excludes the existence of $\\alpha^2$ dynamos in radiative zones and the second property makes the existence of $\\alpha\\Omega$ dynamos (with weak differential rotation in mid-latitudes)  unlikely. As we have shown only a dynamo with  $\\alpha \\propto \\cos\\theta$    produces the toroidal fields in mid-latitudes. For $\\ell>1$, however, the belts are more and more shifted  into the polar region. Such fields   cannot close the loop of field amplification by reproducing the original axisymmetric toroidal field.   Or, with other words, the  rotation law with the considered  profile ($\\Omega\\propto \\cos^2\\theta$) mainly induces toroidal fields at mid-latitudes where for high values of $\\ell$ almost no $\\alpha$ effect exists. Hence, an  $\\alpha\\Omega$ dynamo  could only work for very fast rotation. By very fast rotation the instability is suppressed (see Fig. \\ref{f1}). This topological  problem seems to be the key problem with the magnetic-driven dynamo rather than the  excitation conditions for $\\alpha\\Omega$ dynamos. It is, however,  not yet clear whether this argumentation also holds with the same power with  other than the used rotation laws and/or  in fully nonlinear simulations. All the presented $\\alpha\\Omega$ dynamo models working with a weak   solar-type differential rotation are reproducing the assumed sign of the large-scale current helicity ${\\vec B}\\cdot \\curl {\\vec B}$. If thus the dynamo produces its own $\\alpha$ effect by the magnetic instability then the signs will be consistent.  A basic deficit of the presented theory is the fact that both helicity and $\\alpha$ effect have been computed in rigidly rotating stars while for an $\\alpha\\Omega$ dynamo the rotation must be nonrigid. It might  be the case that growth rates and magnetic patterns of the Tayler instability are strongly modified by even weak differential rotation. The discussion of this new subject, however, is not the scope of the present paper. We have thus considered only  cases with a rather weak  differential rotation (3 \\%).  Another deficit is formed by the order-of-magnitude estimation of the eddy diffusivity. Using the characteristic scales of the modes with the highest growth rates the typical velocity is of  order 1 mm/s and the typical diffusivity value is about $10^5$ cm$^2$/s (both for solar values). Estimations of the radial mixing produced by the same instability provide an important test of the theory. The observed content of light elements at the Sun impose restrictions on radial mixing in the upper radiative core (Barnes et al.  1999, and references therein). The chemical mixing  in the deep stellar interior can also be compatible with observations only if its characteristic time is longer than the evolutionary time scale \\citep{KW94}. Only very slightly supercritical kink-type instability can satisfy this restriction \\citep{KR08}.  All the material  values in the paper are solar values. The formulation of the dynamo theory is such that the material parameters of the stellar interior do not appear (see Eq. (\\ref{AA})). The conclusions about the dynamo activity for hot stars do thus not depend on the mass of the star.  This is not true, however, for the above mentioned  order-of-magnitude estimations of the characteristic values of  velocity and diffusivity which are running with $\\Omega/N$ and $(\\Omega/N)^2$, resp. From Fig. \\ref{f5} one finds that the given solar values increase by  factors of 30 or 900, resp., for stars with three solar masses.   Only  nonlinear simulations can demonstrate whether the system indeed prefers the modes with the highest growth rates. Only if  not then the  radiative-zone dynamo has  a  chance to work in the solar interior  but if it  works then  it should be only marginally supercritical."
    },
    "1107/1107.3314_arXiv.txt": {
        "abstract": "The chemistry of the early Universe is a fascinating field of study. Even in the absence of any elements heavier than lithium, a surprising degree of chemical complexity proves to be possible, giving the topic considerable interest in its own right. In addition, the fact that molecular hydrogen plays a key role in the formation of the first stars and galaxies means that if we want to understand the formation of these objects, we must first develop a good understanding of the chemical evolution of the gas. In this review, I first give a brief introduction to the chemistry occurring in the gas prior to the formation of the first stars and galaxies, and then go on to discuss in more detail the main chemical processes occurring during the gravitational collapse of gas from intergalactic to protostellar densities, and how these processes influence the final outcome of the collapse. ",
        "introduction": "\\label{sec:intro} Chemistry in the early Universe has proved to be an interesting topic of study for several reasons. Before the first stars and galaxies formed, the elemental composition of the gas was very simple -- it consisted only of hydrogen and helium (and their isotopes) plus a tiny trace of lithium -- and yet this is sufficient to give rise to a surprising degree of chemical complexity. In addition, the interaction between the gas and the cosmic microwave background (CMB) plays an important role in regulating the chemistry of the gas, and moreover one which may lead to observable changes in the CMB \\citep[see e.g.][]{maoli94,dub08,sch08}. Furthermore, understanding the chemical evolution of the gas at these early epochs is vital if one wants to understand the formation of the first stars and galaxies, given the crucial role played by molecular hydrogen (H$_{2}$) in this process.  There is a wide range of topics that one could discuss in a review of the chemistry of the early universe, ranging from the physics and chemistry of the recombination epoch at redshift $z > 800$ to the reionization of the Universe at $z < 15$. However, to keep this review to a manageable length, I will focus on only a few key issues, and will not discuss either the recombination epoch or the process of cosmological reionization. Both of these topics are treated in detail elsewhere (see e.g.\\ \\citet{fendt09} and references therein for a lengthy discussion of the physics of the recombination epoch, or the recent review by \\citet{mw10} for a discussion of the epoch of reionization).  In the next Section,  I review what we know about the chemical evolution of the IGM from the end of the recombination epoch until the time at which the earliest protogalaxies form. I follow this up in Section~\\ref{proto} by discussing the chemical evolution of the gas within these early protogalaxies. I show how an understanding of the gas chemistry can allow us to determine the mass scale of the first star-forming minihalos, outline the main chemical changes that occur in the gas during its collapse to protostellar densities, and highlight the important role played by molecular hydrogen throughout this process. ",
        "conclusions": ""
    },
    "1107/1107.1311_arXiv.txt": {
        "abstract": "In cool stars that oscillate like the Sun, non-radial modes become mixed as the stars evolve.  The mixing is caused by the coupling between g-modes in the stellar core and p-modes in the envelope, which results in distinctly different and more complex frequency spectra for subgiants and red giants than seen in main sequence stars.  Using a new version of the `scaled' \\'echelle diagram, I illustrate how the frequencies of non-radial modes evolve during the evolution from the main sequence to the red giant branch, and I show how they depend on stellar mass and metallicity.  Then, with focus on the dipole ($l=1$) modes, which show the strongest effects from mixing, I present a toy model to fit, and hence identify, those modes in a large series red giant models. ",
        "introduction": "Stellar models predict that non-radial oscillation modes in stars like the Sun will become of a mixed character as the stars evolve off the main sequence \\citep{Osaki75}.  The mixing, which is caused by the coupling of two modes (a g-mode in the stellar core and a p-mode in the envelope) will result in shifts of the mode frequencies - known as mode bumping - as the two modes undergo so-called avoided crossings. As a result, mixed modes do not follow the usual asymptotic relation of regular spaced high radial order p-mode frequencies, $\\nu$, as expressed by \\citep[e.g.~][]{Vandakurov68,Tassoul80,Dalsgaard11} \\begin{equation} \\nu_{nl}=\\Delta\\nu(n+l/2+\\epsilon)-\\delta\\nu_{0l} \\label{asymptot} \\end{equation} where the small frequency separation % \\begin{equation} \\delta\\nu_{0l}\\propto(l+1/2)^2\\frac{\\Delta\\nu}{\\nu_{nl}}\\,. \\label{smallsep} \\end{equation} Here $n$ is the radial order, $l$ is the degree, $\\epsilon$ is a dimensionless offset, and \\dnu~is the separation between over tone modes of the same degree. When visualized in the \\'echelle diagram, where we plot mode frequency versus the frequency modulo the large frequency separation, \\dnu, mixed modes can show strong departures from the almost vertical ridges of the non-mixed modes, which I will discuss next. ",
        "conclusions": "I have shown that a new version of the `scaled' \\'echelle diagram can be useful to compare an ensemble of stellar models, and in particular to show the evolution of pulsation frequencies from the main sequence to the RGB as well as how % the frequencies depend on mass and metallicity. For red giants we see that the $l=1$ modes become more confined as small clusters of modes.  While this in some sense makes the frequency spectra simpler, it also brings new challenges by making it harder to determine the frequencies of all the modes between each cluster. I briefly demonstrated that fitting a simple toy model to stellar model frequencies can aid in establishing relations between seismic observables and the toy model parameters. When fully developed this could lead to a robust way to fit the observed $l=1$ frequencies and hence to determine the characteristics of their complex frequency separations."
    },
    "1107/1107.1916_arXiv.txt": {
        "abstract": "If the dark matter consists of supersymmetric particles, $\\gamma$-ray observatories such as the Large Area Telescope aboard the Fermi satellite may detect annihilation radiation from the haloes of galaxies and galaxy clusters. Much recent effort has been devoted to searching for this signal around the Milky Way's dwarf satellites. Using a new suite of high-resolution simulations of galaxy cluster haloes (the Phoenix Project), together with the Aquarius simulations of Milky-Way-like galaxy haloes, we show that higher signal-to-noise and equally clean signals are, in fact, predicted to come from nearby rich galaxy clusters. Most of the cluster emission is produced by small subhaloes with masses less than that of the Sun. The large range of mass scales covered by our two sets of simulations allows us to deduce a physically motivated extrapolation to these small (and unresolved) masses. Since tidal effects destroy subhaloes in the dense inner regions of haloes, most cluster emission is then predicted to come from large radii, implying that the nearest and brightest systems should be much more extended than Fermi's angular resolution limit. The most promising targets for detection are clusters such as Coma and Fornax, but detection algorithms must be tuned to the predicted profile of the emission if they are to maximize the chance of finding this weak signal. ",
        "introduction": "Annihilation radiation at $\\gamma$-ray frequencies offers one of the most exciting prospects for non-gravitational detection of cold dark matter, and is expected if the dark matter consists of supersymmetric particles~\\citep[e.g.][]{Berezinsky1994,Berezinsky2003,Bergstrom1998,Stoehr2003,Koushiappas2004, Colafrancesco2007,Diemand2007,Kuhlen2008,Pieri2008,sp08a,Strigari2008,jeltema10,ackermann10,Zavala2010}. Much effort is being devoted to searching for this signal around the Milky Way's dwarf companions, in particular using the Fermi satellite~\\citep{abdo}.  Predictions for the properties of the annihilation radiation rely on a detailed understanding of the structure of cold dark matter haloes which can be gained only through high-resolution numerical simulations of halo formation.  The structure of galaxy-mass cold dark matter haloes has been investigated in considerable depth~\\citep[e.g.][]{Diemand2007,diemand08,Kuhlen2008,sp08a,sp08b,Anderson2010,Kamionkowski2010} showing that the radial distribution of low-mass subhaloes, and thus of annihilation radiation, is much less centrally concentrated than that of the dark matter as a whole. In the Milky Way, this results in the dominant subhalo contribution to the annihilation radiation coming from large galactrocentric distance and so appearing almost uniform across the sky to an observer on Earth~\\citep{sp08a}. This same effect causes the annihilation radiation from an external galaxy cluster to appear much less centrally concentrated than the distribution of galaxies. As we show below, this has significant implications for the optimal strategy for detecting the annihilation signal.  In this paper we present some of the largest high-resolution simulations of cluster haloes to date (the Phoenix Project) and use them to investigate the detailed structure of the dark matter distribution in clusters and its halo-to-halo variation. We use these data, together with data from the Aquarius set of galaxy halo simulations \\citep{sp08b}, to predict the expected $\\gamma$-ray annihilation radiation from cluster haloes which we compare to the expected annihilation radiation from giant and satellite galaxy haloes.  As we were completing this work, \\cite{Pinzke11} and \\cite{sanchez} posted preprints investigating, amongst other things, the $\\gamma$-ray annihilation radiation expected from galaxy clusters. The luminosity and spatial distribution of this radiation depend sensitively on the properties of surviving dark matter subhaloes down to the limiting mass of the cold dark matter power spectrum, which may be in the range $10^{-6}$ to $10^{-12}M_\\odot$~\\citep{Hofmann2001,Green05}.  For their analysis, \\cite{Pinzke11} relied on an extrapolation of scalings based on published results for simulations of galactic dark matter haloes, including those of the Aquarius Project, while \\cite{sanchez} extended the semi-analytic model of \\cite{Kamionkowski2010}, rescaling relevant model parameters. Combining the Phoenix and Aquarius simulations we test explicitly the validity of the scalings used by \\cite{Pinzke11} and \\cite{sanchez}, we investigate their underlying physical basis, and we thus construct a more robust (though still uncertain) framework for extrapolation. For the most part, our results are in agreement with those of \\cite{Pinzke11}, but not with those of \\cite{sanchez} who infer a much weaker contribution from subhalos to the total annihilation radiation from clusters than \\cite{Pinzke11} or us find. In this study, we also present an estimate of the expected signal-to-noise of the annihilation radiation from nearby clusters and compare it to that from nearby dwarf and giant galaxies.  The outline of our paper is as follows. Sections~2 and~3 give brief descriptions of our simulation suite and of a model for calculating the annihilation flux and its signal-to-noise in an idealised experiment. In Section~4, we discuss our results and their implications for dark matter detection. ",
        "conclusions": "\\begin{table*} \\begin{center} \\begin{tabular}{lcccccc} \\hline \\hline Object Name & Half-light radius & Distance & $M_{\\rm 200}$  & $L$ &$F=L/(4\\pi d^2)$ & $S/N$ \\\\  &[arcmin] &[Mpc] &[${\\rm M_\\odot}$] &[$L_{\\rm mw}$] &[$F_{\\rm mw}$] &[$(S/N)_{\\rm mw}$]\\\\ \\hline \\hline AWM 7  & 35.5 & 67.0 & $4.2 \\times 10^{14}$ &$7.1 \\times 10^4$ & $3.2 \\times 10^{-4}$ & $6.8 \\times 10^{-3}$ \\\\ Fornax Cluster & 84.1 & 17.5 & $1.0 \\times 10^{14}$ &$1.2 \\times 10^4$ & $8.0 \\times 10^{-4}$ & $7.3 \\times 10^{-3} $\\\\ M49    & 59.6 & 18.2 &$0.4 \\times 10^{14}$ &$3.9 \\times 10^3$ &$2.4 \\times 10^{-4}$ &$3.1 \\times 10^{-3}$ \\\\ NGC 4636 &52.6 &  17.4 & $0.24 \\times 10^{14}$ &$2.1 \\times 10^3$ &$1.4 \\times 10^{-4}$ &$2.0 \\times 10^{-3}$ \\\\ Centaurus (A3526) &40.1  & 50.5 & $2.6 \\times 10^{14}$ &$3.9 \\times 10^4$ &$3.1 \\times 10^{-4}$ &$5.8 \\times 10^{-3}$ \\\\ Coma & 36.1 & 95.8 &$1.3 \\times 10^{15}$ &$2.9 \\times 10^5$ &$6.4 \\times 10^{-4}$ &$1.3 \\times 10^{-2}$ \\\\ \\hline \\hline Draco &16.4 &0.082 & N/A &$5.2 \\times 10^{-3}$ &$1.6 \\times 10^{-5}$ &$6.3 \\times 10^{-4}$ \\\\ UMaI  &18.4  &0.066 & N/A &$4.3 \\times 10^{-3}$ &$2.0 \\times 10^{-5}$ & $7.5 \\times 10^{-4}$ \\\\ LeoI &4.4 &0.25 & N/A &$3.5 \\times 10^{-3}$ & $1.2 \\times 10^{-6}$ &$8.2 \\times 10^{-5}$ \\\\ Fornax dwarf &5.9 & 0.138 & N/A &$2.0 \\times 10^{-3}$ &$2.2 \\times 10^{-6}$ &$1.5 \\times 10^{-4}$ \\\\ LeoII &2.5 & 0.205 & N/A &$8.5 \\times 10^{-4}$ &$4.1 \\times 10^{-7}$ &$3.1 \\times 10^{-5}$ \\\\ Carina &4.6 & 0.101 & N/A &$7.1 \\times 10^{-4}$ &$1.4 \\times 10^{-6}$ &$1.0 \\times 10^{-4}$ \\\\ Sculpt &13.2 & 0.079 & N/A &$3.2 \\times 10^{-3}$ &$1.0 \\times 10^{-5}$ &$4.9 \\times 10^{-4}$ \\\\ Sext &3.3 & 0.086 & N/A &$3.0 \\times 10^{-4}$ &$8.3 \\times 10^{-7}$ &$6.1 \\times 10^{-5}$ \\\\ UMaII &28.8 &0.032 & N/A &$2.6 \\times 10^{-3}$ &$5.2 \\times 10^{-5}$ &$1.3 \\times 10^{-3}$ \\\\ Comber &15.9 &0.044 & N/A &$1.6 \\times 10^{-4}$ &$1.7 \\times 10^{-5}$ &$6.8 \\times 10^{-4}$ \\\\ WilI &17.7 &0.066 & N/A &$3.9 \\times 10^{-3}$ &$1.8 \\times 10^{-5}$ &$7.0 \\times 10^{-4}$ \\\\ LMC &82.5 &0.049 & N/A &$3.8 \\times 10^{-2}$ &$3.3 \\times 10^{-4}$ &$3.1 \\times 10^{-3}$ \\\\ SMC &45.5 &0.061 & N/A &$1.9 \\times 10^{-2}$ &$1.1 \\times 10^{-4}$ &$1.8 \\times 10^{-3}$ \\\\ \\hline \\hline M31 &351.5 &0.807 &$1.8 \\times 10^{12}$ &$1.3 \\times 10^{2}$ &$4.2 \\times 10^{-3}$ &$9.3 \\times 10^{-3}$ \\\\ \\hline \\hline \\end{tabular} \\end{center} \\caption{Principal properties of nearby galaxy clusters, prominent satellites of the Milky Way, and the Andromeda Nebula, M31. The annihilation luminosity, $L$, is given in units of the luminosity from the smooth component of the main Aq-A halo, which we use as a proxy for the Milky Way.  The observed flux, $F$, is expressed relative to the flux received by an observer placed $8\\,{\\rm kpc}$ from the centre of Aq-A. Similarly, the predicted $S/N$ for an optimal filter placed on each object is normalized to the signal-to-noise predicted for a similar filter tuned to the diffuse emission of Aq-A seen from this observer location. For the signal-to-noise calculations, we use the optimal filter of~Springel et al. (2008a) assuming the background to be the same everywhere and to dominate the signal in all objects.} \\label{TabObes} \\end{table*}   In Table~\\ref{TabObes} we summarise properties of some nearby astronomical objects which are relevant for the detectability of their dark matter annihilation signal.  We consider six galaxy clusters which were already analyzed by the Fermi collaboration \\citep{ackermann10}, thirteen of the known dwarf satellites of our Galaxy, and the nearest giant galaxy, M31.  For the galaxy clusters, distances were taken from the NASA/IPAC Extragalactic Database\\footnote{http://nedwww.jpac.caltech.edu/}, and virial masses, ${\\rm M_{200}}$ (based on X-ray data), from \\cite{reiprich02}. Values for $V_{\\rm max}$ and $r_{\\rm max}$ were derived assuming an NFW density profile~\\citep{NFW96,NFW97} and the mass-concentration relation of \\cite{Neto2007}. We have verified that this relation is consistent with our simulation data down to the resolution limit of Aq-A-1, which is about $10^5\\, {\\rm M_{\\odot}}$.  Data for dwarf satellites were taken from the mass models of \\cite{penarrubia08}. Their $\\gamma$-ray luminosities are estimated from an emission integral based on the NFW formula, $\\int \\rho^2 dV = 1.23\\,V_{\\rm max}^4/(G^2r_{\\rm max})$. As discussed in \\cite{sp08a}, the annihilation signal due to substructures within Milky Way dwarfs (the `subsub' component) is less than that due to the smooth component of their haloes in almost all cases, so we do not consider it here. The distance of M31 was also taken from the NASA/IPAC Extragalactic Database. We base structural parameters for the M31 halo on the Aq-A-1 simulation which has a very similar mass~\\citep{lw08}.  We estimate a ``best case'' signal-to-noise for each object using the optimal filter discussed by \\cite{sp08a} and assuming a uniform background accross the whole sky which dominates over the signal in all objects.  In this case, the optimal filter has the same shape as the signal, and the signal-to-noise can be written in the generic form \\begin{equation} S/N = f_{\\rm shape}(\\theta_h/\\theta_{\\rm psf})\\,\\left[\\frac{t A_{\\rm eff}}{B}\\right]^{1/2} \\frac{F}{(\\theta^2_{h}+\\theta^2_{\\rm psf})^{1/2}}, \\end{equation} where $F = L/(4\\pi d^2)$ is the photon flux, $\\theta_h$ the half-light radius, $\\theta_{\\rm psf}$ ($\\simeq 10$ arcmin for Fermi at the relevant energies~\\citep{Michelson07}) describes the point spread function of the instrument, $t$ is the integration time, $A_{\\rm eff}$ is the effective collecting area of the telescope, and $B$ is the background count rate per unit solid angle. The function $f_{\\rm shape}(x)$ encodes the detailed shape of the emission profile of the signal~\\citep{sp08a}; it is of order unity and depends only weakly on the ratio $x=\\theta_h/\\theta_{\\rm psf}$.  Using the techniques discussed above, we can compare the apparent $\\gamma$-ray luminosities and achievable $S/N$ ratios for galaxy clusters with those estimated by \\cite{sp08a} for dwarf satellites of the Milky Way. Results are shown in the Table. We find that the brightest nearby cluster, Fornax, is predicted to appear 15 times more luminous than the brightest dwarf spheroidal, UMaII, and 40-50 times more luminous than UMaI, Draco or the ultrafaint satellite, Wilman-1. However, the Fornax cluster is quite extended on the sky, and, as a result, when optimal filters are used, the slightly fainter but more compact Coma cluster has a predicted $S/N$ ratio 1.8 times larger and ten times that of the most easily detectable dwarf spheroidal, UMaII. Although the Andromeda Nebula is predicted to have comparable $S/N$, it is not a promising target because of the difficulty in correcting for foreground and other sources of emission. Note that the $S/N$ predicted for both objects is still very small compared to that of the main component of the Milky Way's smooth halo. Here also, of course, the main problem is in separating annihilation radiation from other $\\gamma$-ray signals.  The Coma cluster thus offers an order of magnitude better opportunity than any Milky Way satellite for detecting dark matter or placing limits on its annihilation cross-section. As we have shown, for a high resolution experiment like Fermi, the sensitivity for detecting such radiation will be enhanced by use of a filter which is properly matched to the expected extent of the object. For example, for the optimal filter, the $S/N$ expected for Coma is about 1.5 times higher than the $S/N$ for a filter based on the point-spread function of the Fermi LAT, assuming 10 arcmin for the latter at the relevant energies. Detecting annihilation radiation from the Coma or Fornax clusters or the placing of robust and stringent upper limits will also require careful subtraction of astrophysical sources and an accurate estimate of the background."
    },
    "1107/1107.3913_arXiv.txt": {
        "abstract": "The Red MSX Source (RMS) survey has identified a sample of $\\sim$1200 massive young stellar objects (MYSOs), compact and ultra compact H{\\sc ii} regions from a sample of $\\sim$2000 MSX and 2MASS colour selected sources. We have used the 100-m Green Bank telescope to search for 22-24\\,GHz water maser and ammonia (1,1), (2,2) and (3,3) emission towards $\\sim$600 RMS sources located within the northern Galactic plane. We have identified 308 H$_2$O masers which corresponds to an overall detection rate of $\\sim$50\\,per\\,cent. We find no significant difference in the detection rate for H{\\sc ii} regions and MYSOs which would suggest that the conditions required to produce maser emission are equally likely in both phases. Comparing the detection rates as a function of luminosity we find the H$_2$O detection rate has a positive dependence on the source luminosity, with the detection rate increasing with increasing luminosity.  We detect ammonia emission towards 479 of these massive young stars, which corresponds to $\\sim$80\\,per\\,cent. Ammonia is an excellent probe of high density gas allowing us to measure key parameters such as gas temperatures, opacities, and column densities, as well as providing an insight into the gas kinematics. The average kinetic temperature, FWHM line width and total NH$_3$ column density for the sample are approximately 22\\,K, 2\\,\\kms\\ and $2\\times 10^{15}$\\,cm$^{-2}$, respectively. We find that the NH$_3$ (1,1) line width and kinetic temperature are correlated with luminosity and finding no underlying dependence of these parameters on the evolutionary phase of the embedded sources, we conclude that the observed trends in the derived parameters are more likely to be due to the energy output of the central source and/or the line width-clump mass relationship.  The velocities of the peak H$_2$O masers and the NH$_3$ emission are in excellent agreement with each other, which would strongly suggest an association between the dense gas and the maser emission. Moreover, we find the bolometric luminosity of the embedded source and the isotropic luminosity of the H$_2$O maser are also correlated. We conclude from the correlations of the cloud and water maser velocities and the bolometric and maser luminosity that there is a strong dynamical relationship between the embedded young massive star and the H$_2$O maser. ",
        "introduction": "Massive young stellar objects (MYSOs) and ultra compact (UC) H{\\sc ii} regions are two of the earliest phases in the lives of OB stars. The MYSO phase begins when heating of the proto stellar envelope increases visibility of the object in the mid-infrared, and ends once the central star begins to ionize its surrounding environment and forms an UC\\,H{\\sc ii} region. These two phases are physically distinct from the earlier hot molecular core phase, which is not generally detectable at mid-infrared wavelengths (e.g., \\citealt{de-buizer2002}). MYSOs also possess strong ionized stellar winds (e.g., \\citealt{bunn1995}); however, the radio emission from these winds is relatively weak ($\\sim$1\\,mJy at 1\\,kpc; \\citealt{hoare1994}) and easily distinguishable from the slightly later radio-loud UC\\,H{\\sc ii} region phase. It is likely that many MYSOs and UC\\,H{\\sc ii} regions are still accreting, as evidenced by their almost ubiquitous association with powerful bipolar outflows (\\citealt{lada1985}).  Models suggest that high levels of accretion in the early stages of a massive star's development lead to a `swelling up' of the protostellar core due to trapped entropy (e.g., \\citealt{yorke2008,hosokawa2009,hosokawa2010}). The length of this swollen phase is comparable to the Kelvin-Helmholz timescale, which for stars with a current mass of $\\ga$20\\,\\msun\\ is only a few thousand years. This means that once the accreted mass exceeds $\\sim$20\\,\\msun\\ the star arrives on the main-sequence, even if accretion is still ongoing \\citep[see also][]{mckee2003}. At this point, the star begins to ionize its surroundings and becomes a UC\\,H{\\sc ii} region. This means that the MYSO phase is likely to be very brief ($\\sim 10^{5}$\\,yrs) and is limited to objects with {\\it current} masses lower than $\\sim$20\\,\\msun\\ (\\citealt{davies2011}). The UC\\,H{\\sc ii} region phase lasts a factor of 2-5 times longer, with a small dependence on the {\\it final} stellar mass (\\citealt{davies2011,mottram2011b}).   The Red MSX Source (RMS; \\citealt{hoare2005}; \\citealt{mottram2006}; \\citealt{urquhart2007c}) Survey has established a large ($\\sim$1200)  and well-selected sample of MYSOs and compact and UC\\,H{\\sc ii} regions, and a database of complementary multi-wavelength data.\\footnote{www.ast.leeds.ac.uk/cgi-bin/RMS/RMS\\_DATABASE.cgi.} We have used arcsecond-resolution mid-infrared imaging from the Spitzer GLIMPSE survey (\\citealt{benjamin2003}) or our own ground-based imaging (e.g., \\citealt{mottram2007}) to reveal multiple and/or extended sources within the MSX beam, as well as MYSOs in close proximity to existing H{\\sc ii} regions. We have obtained arcsecond-resolution radio continuum with ATCA and the VLA (\\citealt{urquhart_radio_south,urquhart_radio_north}) to identify UC\\,H{\\sc ii} regions and PNe, whilst observations of $^{13}$CO transitions (\\citealt{urquhart_13co_south,urquhart_13co_north}) deliver kinematic distances and luminosities, which allow us to distinguish between nearby low- and intermediate-mass YSOs and genuine MYSOs. Finally we have obtained near-infrared spectroscopy (e.g., \\citealt{clarke2006}) which allows us to distinguish the more evolved stars.   With the source classification effectively complete, the next step is to examine the global characteristics of this Galaxy-wide sample of massive young stars. This involves determining the physical and chemical nature of the environment as a way of gauging the evolutionary status of our sample of MYSOs and H{\\sc ii} regions. In this paper we present the results of a set of thermal ammonia (NH$_3$) and water (H$_2$O) maser observations made with the 100-m Green Bank telescope (GBT) towards a sample of $\\sim$600 young massive stars.  Ammonia is an excellent tracer of high-density gas ($\\sim$10$^4$\\,cm$^{-3}$; \\citealt{evans1999,stahler2005}), and is relatively unaffected by depletion at lower temperatures compared to other common molecular tracers such as CO (\\citealt{bergin1997}). The NH$_3$ (1,1), (2,2) and (3,3) inversion transitions are normally collisionally excited and the rotational temperature of the gas can be determined from the intensity ratio of any two inversion transitions. The inversion transition splits into 18 separate hyperfine components (\\citealt{ho1983}), usually resolved into five distinct components, the ratio of which can be used to calculate the optical depth of the transition.  H$_2$O masers are known to occur in both high and low-mass star-forming regions (e.g., \\citealt{forster1999,claussen1996}) and are thus an important signpost of ongoing star formation. Water masers are generally thought to be associated with molecular outflows (\\citealt{codella2004} and references therein). We will use the H$_2$O detection rates for our sample of MYSOs and H{\\sc ii} regions to investigate their statistical association as a function of evolutionary phase and luminosity.  The structure of the paper is as follows: in Sect.\\,2 we discuss the source selection, the observational setup and the data reduction procedure. In Sect.\\,3 we use the reduced spectra to determine physical properties of the star-forming environments. In Sect.\\,4 we present the results and discuss their implications. We present a summary of our results and highlight our main findings in Sect.\\,5. ",
        "conclusions": "\\label{sect:summary}  We have used the 100-m Green Bank Telescope to conduct a programme of 22-24\\,GHz spectral line observations towards 597 massive young stellar objects (MYSOs) and compact and ultra-compact H{\\sc ii} regions identified by the Red MSX Source (RMS) survey. We targeted the 22\\,GHz water (H$_2$O) maser transition and thermal ammonia (NH$_3$) (1,1), (2,2) and (3,3) inversion transitions.  These observations have an angular resolution of $\\sim$30\\arcsec\\ and typical 4$\\sigma$ sensitivity of $T_{\\rm{mb}}= 0.2$\\,K per 0.32-km\\,s$^{-1}$ channel.  We detect water maser emission towards 308 RMS sources with an overall detection rate of $\\sim$50\\,per\\,cent. NH$_3$ emission is detected towards 479 sources ($\\sim$80\\,per\\,cent of the sample) with the NH$_3$ (1,1) hyperfine structure being clearly detected towards $\\sim$400 sources, and thus allowing the kinetic temperature, optical depth and column density to be determined for a large number of massive star-forming regions. The average kinetic temperature, FWHM line width and column density for the sample are approximately 22\\,K, 2\\,\\kms\\ and $2\\times 10^{15}$\\,cm$^{-2}$, respectively.  We find the detection rates for both the H$_2$O masers and NH$_3$ emission is similar for both MYSO and H{\\sc ii} region sub-samples. Combining these results with a large database of complementary observations we analyse these parameters as a function of luminosity and evolutionary phase, and correlate these physical conditions with H$_2$O maser emission. Our main findings are as follows:   \\begin{enumerate}  \\item There is no significant difference in the H$_2$O maser detection rate for H{\\sc ii} regions and MYSOs which would suggest that the conditions required to produce maser emission are equally likely in both phases. Comparing the detection rates as a function of luminosity we find the H$_2$O detection rate has a positive dependence on the source luminosity, with the detection rate increasing with increasing luminosity.   \\item We find that the NH$_3$ (1,1) line width and kinetic temperature are correlated with luminosity. It is clear from these correlations that the embedded sources are having measurable effects on their surrounding molecular environments. Finding no underlying dependence of these parameters on the evolutionary phase of the embedded sources, we conclude that the observed trends in the derived parameters are related to the energy output of the central source and/or the linewidth-clump mass relationship.  \\item The velocities of the peak H$_2$O masers and the NH$_3$ emission for each relevant source are in excellent agreement with each other with a correlation coefficient of 0.92, which would strongly suggest an association between the dense gas and the maser emission. Moreover, we find the bolometric luminosity of the embedded source and the isotropic luminosity of the H$_2$O maser are also correlated. We conclude that there is a strong dynamical relationship between the embedded young massive star and the H$_2$O maser.  \\item We find excellent agreement between velocities of sources using $^{13}$CO observations and those obtained from NH$_3$ transitions. We find only three sources where the velocity assigned using the CO data is incorrect.  \\end{enumerate}  These observations are a first step in examining the global characteristics of this Galaxy-wide sample of massive young stars and will form the cornerstone for more detailed studies of well selected sub-samples."
    },
    "1107/1107.4857.txt": {
        "abstract": "Improvements of in-orbit calibration of GSO scintillators in the Hard X-ray Detector on board Suzaku are reported. To resolve an apparent change of the energy scale of GSO which appeared across the launch for unknown reasons, consistent and thorough re-analyses of both pre-launch and in-orbit data have been performed. With laboratory experiments using spare hardware, the pulse height offset, corresponding to zero energy input, was found to change by $\\sim 0.5$\\% of the full analog voltage scale, depending on the power supply. Furthermore, by carefully calculating all the light outputs of secondaries from activation lines used in the in-orbit gain determination, their energy deposits in GSO were found to be effectively lower, by several percent, than their nominal energies. Taking both these effects into account, %the in-orbit data were confirmed to be consistent with the on-ground measurements. the in-orbit data agrees with the on-ground measurements within $\\sim $5\\%, without employing the artificial correction introduced in the previous work (Kokubun et al. 2007). %With this knowledge, the GSO data processing and the response matrix of GSO were updated. %As a result, the HXD-PIN and HXD-GSO spectra of the Crab Nebula are %now expressed successfully over 12-300 keV by a broken powerlaw %with a break energy of $\\sim 110$ keV. With this knowledge, we updated the data processing, the response, and the auxiliary files of GSO, and reproduced the HXD-PIN and HXD-GSO spectra of the Crab Nebula over 12-300 keV by a broken powerlaw with a break energy of $\\sim 110$ keV. ",
        "introduction": "\\label{section:1}  %\\green{\u0082\u00d0\u0082\u00b6\u0082\u00e5\u0082\u00a4\u0082\u00c9\u008a\u00c8\u0092P\u0082\u00c8\u0082\u00b7\u0082\u00b4\u0082\u00ad\u0082\u00c6HXD\u008f\u00d0\u0089\u00ee} The fifth Japanese X-ray satellite Suzaku (\\cite{Mitsuda2007}), carries the X-ray Imaging Spectrometer (XIS; \\cite{Koyama2007}) located at the foci of the X-ray Telescope (XRT \\cite{Serlemitsos2006}) and a non-imaging hard X-ray instrument, the Hard X-ray Detector (HXD). The HXD covers a hard X-ray energy range of 10--600 keV, utilizing Si-PIN diodes (hereafter HXD-PIN), and gadolinium silicate scintillators (Gd$_2$SiO$_5$:Ce, hereafter HXD-GSO) which are placed behind HXD-PIN. The detailed design of the HXD is summarized in Takahashi et al. (2007) (hereafter Paper I), followed by a report on its in-orbit performance (\\cite{Kokubun2007}; hereafter Paper II) and that of the timing property \\citep{Terada2008}.  %\\green{\u0092n\u008f\u00e3\u0082\u00c5\u0082\u00e0\u008f\u00e3\u008b\u00f3\u0082\u00c5\u0082\u00e0cal\u0082\u00b5\u0082\u00bd} Prior to the launch on 2005 July 10, the performance of the HXD was tested and calibrated at every stage of its integration, and also after the instrument was mounted on the spacecraft. In addition to these pre-launch experiments, several Monte-Carlo simulations were performed, to better understand the HXD performance (including its background) expected in near-Earth radiation environments (\\cite{Terada2005}; Paper I). We have also calibrated the HXD performance extensively after the launch (Paper II), and compared all propertied that can be obtained in-orbit with the results of the ground measurements. As a result, we have successfully verified that the HXD performs in-orbit generally as expected (Paper II), and reconfirmed most of the results of the pre-launch calibration.  %\\green{\u0082\u00bb\u0082\u00cc\u008c\u008b\u0089\u00ca\u0081A\u0096\u00b5\u008f\u0082\u0082\u00aa\u008ec\u0082\u00c1\u0082\u00c4\u0082\u00a8\u0082\u00e8\u0081A\u0082\u00bb\u0082\u00cc\u0089\u00f0\u008c\u0088\u0082\u00aa\u0096\u00da\u0093I}  In spite of these successful in-orbit calibration efforts, we are still left with a few unsolved issues of the HXD performance. One of them is an apparent change in the energy scale (relation between the incident X-ray energy and detector pulse height) of HXD-GSO. We found this issue when we observed the X-ray pulsar A0535+26, which exhibits an absorption line at $\\sim 50$ keV \\citep{Terada2006}; this spectral feature appeared at discrepant energies in the HXD-PIN and HXD-GSO spectra that were produced based on the nominal pre-launch calibration (see subsection 2.3). Since the energy scale of HXD-PIN was kept unchanged across the launch (figure 5 of Paper II), we tentatively concluded that the energy scale of HXD-GSO changed for some unspecified reasons, and hence added an artificial non-linearity to the GSO energy scale (figure 12 of Paper II) as a temporary solution.  The aim of the present paper is to more fundamentally solve this issue, by revisiting the pre-launch and in-orbit calibrations, and conduct some laboratory experiments using spare hardware. Section 2 is used to review the way of conversion from energy to pulse height, and to present the investigation strategy. In section 3, we study the spare hardware performance, followed in section 4 by a detailed comparison between the on-ground and in-orbit data, and construction of a revised energy scale. We present in section 5 how well our new solution worked on actual observation data, and summarized these works in section 6.  %================== % S.2 % ",
        "conclusions": "We have performed re-calibration of the energy scale of the GSO scintillators which covers the harder energy ranges of the Suzaku HXD. Through laboratory experiments using flight-spare hardware, we found that the pedestal of HXD-AE decreased by 8 ch when the HXD is driven by the spacecraft power supply, than the case where a more stable laboratory power supply is used. Calculating all light outputs of secondaries from activation lines revealed that the energy of calibrators, used to determine the in-flight gain of GSO, was previously over-estimated by several percent. Taking both effects into account, the in-orbit data have been confirmed to be consistent with the pre-launch calibration.  When the GSO energy scale was thus revised, the in-orbit GSO branch width became also consistent with the ground one. Taking into account the energy scale thus updated, as well as the GSO branch width, the response matrixes have been created. Then, the Crab Nebula spectra of HXD-PIN and HXD-GSO, over 12--300 keV, has been expressed by a broken powerlaw with a break energy of $\\sim 110$ keV and a photon index of $\\sim$ 2.1 below this energy.  As a result of these works, major issues with the in-orbit calibration of HXD-GSO have been successfully solved.  \\vspace{0.2cm}  The authors would like to express their thanks to all who have contributed to the design, development, operation, and the calibration of the HXD experiment, and all of the Suzaku supporters. The present work was supported by Grant-in-Aid for JSPS Fellows.   \\appendix"
    },
    "1107/1107.5585.txt": {
        "abstract": "In recent years, accurate observational constraints become available for an increasing number of Galactic X-ray binaries. Together with proper motion measurements, we could reconstruct the full evolutionary history of X-ray binaries back to the time of compact object formation. In this paper, we present the first study of the persistent X-ray source Cygnus\\;X-1 that takes into account of all available observational constraints. Our analysis accounts for three evolutionary phases: orbital evolution and motion through the Galactic potential after the formation of black hole (BH), and binary orbital dynamics at the time of core collapse. We find that the mass of the BH immediate progenitor is $15.0 - 20.0$ M$_\\sun$, and at the time of core collapse, the BH has potentially received a small kick velocity of $\\le 77$ km s$^{-1}$ at 95\\% confidence. If the BH progenitor mass is less than $\\sim 17$ M$_\\sun$, a non zero natal kick velocity is required to explain the currently observed properties of Cygnus\\;X-1. Since the BH has only accreted mass from its companion's stellar wind, the negligible amount of accreted mass is impossible to explain the observationally inferred BH spin of $a_* > 0.95$, and the origin of this extreme BH spin must be connected to the BH formation itself. Right after the BH formation, we find that the BH companion is a $19.8 - 22.6$ M$_\\sun$ main sequence star, orbiting the BH at a period of $4.7 - 5.2$ days. Furthermore, recent observations show that the BH companion is currently super-synchronized. This super-synchronism indicates that the strength of tides exerted on the BH companion should be weaker by a factor of at least two compared to the usually adopted strength. ",
        "introduction": "In recent years, the number of observed black hole (BH) X-ray binaries (XRBs) has grown significantly. For these binaries, there exists a wealth of observation information about their current physical state:  BH and donor masses, orbital period, donor's position on the H-R diagram and surface chemical composition, transient or persistent X-ray emission, and Roche lobe overflow (RLO) or wind-driven character of the mass transfer (MT) process. Furthermore, proper motions have been measured for a handful of these binaries \\citep[e.g.][]{Mirabel...01, Mirabel...02, Mirabel-Rodrigues03}. Together with the earlier measurements of center-of-mass radial velocities and distances, we can obtain information about the three-dimensional kinematic properties of these binaries. Given this plethora of observation results, the current observed sample of BH XRBs provides us with a unique opportunity to understand the formation and evolution of BHs in binaries. This paper is the third in a series where we investigate in detail the BH formation in XRBs, especially focusing on the mass relationship between BHs and their immediate progenitors and the possible BH natal kick magnitude imparted during the core collapse event.  In the first paper of this series, \\citet[][hereafter Paper I]{Bart...05} showed how using the currently available constraints one could uncover the evolution history of an XRB from the present state back to the time just prior to the core collapse event. They applied their analysis to the BH XRB GRO J1655-40. In the second paper, \\citet[][hereafter Paper II]{Tassos...09} performed the same analysis for the case of the BH XRB XTE J1118+480. In this work, we focus on the case of the BH XRB Cygnus X-1. The mass transfer mechanism in Cygnus X-1 is different from the XRBs studied in previous cases. Both donors in GRO J1655-40 and XTE J1118+480 are transferring mass to the BH under Roche lobe overflow, whereas the BH in Cygnus X-1 is accreting mass from the stellar wind of its companion.  The plan of the paper is as follows. In Section 2, we review Cygnus X-1's currently available observational constraints. A general outline of the analysis used to reconstruct the system's evolutionary history is presented in Section 3, while individual steps of the analysis are discussed in more detail in Section $4 - 7$. In Section 8, we derive constraints on the formation of the BH. The final section is devoted to a summary of our results and discussion of some of the assumptions introduced in our analysis. ",
        "conclusions": "In this paper we constrained the progenitor properties and the formation of the BH in the persistent XRB Cygnus\\;X-1. Our analysis accounts for the orbital evolution and motion through the Galactic potential right after the BH formation, and the binary orbital dynamics at the time of core collapse. We find that the mass of the BH immediate progenitor falls within a range of $15.0 - 20.0$ M$_\\sun$ at 95\\% confidence. We note that the maximum progenitor mass is constrained by our adopted upper limit, which is discussed in Section\\;9.1. The BH has potentially received a small natal kick velocity of $\\le 77$ km s$^{-1}$ at 95\\% confidence. In fact if the progenitor mass is less than $\\sim 17$ M$_\\sun$, a non zero natal kick velocity is \\textit{necessary} to explain the currently observed properties of Cygnus\\;X-1. Since the BH has only accreted mass from its companion's stellar wind, the total amount of mass accreted since the BH formation is less than $\\sim 2 \\times 10^{-3}$ M$_\\sun$. This indicates that the observationally inferred BH spin of $a_* > 0.95$ \\citep{gou...11} cannot be explained by mass accretion and has to be natal. This high spin has implications about BH formation and the role of rotation in core collapse. Right after the BH formation, the BH companion has a mass of $19.8 - 22.6$ M$_\\sun$, in an orbit with period of $4.7 - 5.2$ days and eccentricity of $0.015 - 0.022$. Although the post-SN orbital eccentricity is small, the pre-SN orbit can potentially be fairly eccentric. This is possible if the BH receives a natal kick velocity at the right magnitude and direction.  The formation of the BH in Cygnus\\;X-1 has been previously studied by \\cite{nelemans...99} and \\cite{Mirabel-Rodrigues03}. Both studies assumed symmetric mass loss during the core collapse event, and considered only the binary orbital dynamics at the time of core collapse. Comparing with these two earlier studies, we consider the possible asymmetries developed during the core collapse event and the evolution of the binary since the BH formation. It is important to note that these two earlier studies do not consider the multitude of the observational constraints taken into account here and hence the suggested progenitors are not complete solutions for the evolutionary history of Cygnus\\;X-1.  Finally, we discuss some of the assumptions introduced in our analysis in the following sub-sections.  \\subsection{Maximum BH Progenitor Mass}  Unlike the case of GRO J1655-40 studied in Paper I, the analysis of orbital dynamics during the core collapse event does not give an upper limit on M$_{\\rm He}$. Instead, we have conservatively adopted an upper limit of M$_{\\rm He}$ $\\le$20 M$_\\sun$, based on physics involved in the evolution of massive stars. As mentioned in Section 6.1 of Paper II, by evolving a ZAMS star of $\\sim$100 M$_\\sun$ at solar metallicity, the maximum Helium star mass one can achieve is $\\sim$15 M$_\\sun$ when including moderate stellar rotation, and $\\sim$17.5 M$_\\sun$ when assuming no stellar rotation. When adopting the upper limit of 17.5 M$_\\sun$, the lower limit of M$\\rm_{He}$ decreases slightly to 14.6 M$_\\sun$ and the range of V$\\rm_{k}$ becomes $14 - 81$ km\\;s$^{-1}$, both with 95\\% confidence. This range of V$\\rm_k$ still suggests that the BH in Cygnus\\;X-1 received a low kick during the core collapse event.  \\subsection{Association with Cyg OB3}  The center of Cyg OB3 locates at $l = 72.8^\\circ$ and $b = 2.0^\\circ$, and at a distance of $1.4 - 2.7$ kpc away from the Sun \\citep{massey...95, dambis...01, melnik...01, melnik-dambis09}. When comparing that to the location of Cygnus\\;X-1 (Table\\;1), it is clear that not only their Galactic coordinates are close to each other, but also their distance estimations overlap with each other. Furthermore, the measurements of proper motion and radial velocity show that Cygnus X-1 is moving as the members of Cyg OB3 \\citep{dambis...01, Mirabel-Rodrigues03, melnik-dambis09}. Based on these observations, \\cite{Mirabel-Rodrigues03} argue that Cyg OB3 is the parent association of Cygnus X-1. This infers that V$_{\\rm pec,postSN}$ due to the core collapse event has to be small. If we change the constraint on V$_{\\rm pec,postSN}$ to $\\le 10$ km s$^{-1}$, we find M$_{\\rm He}$ to be in a range of $13.9 - 16.9$ M$_\\sun$ and V$_{\\rm k}$ to be $\\le 24$ km/s, both at 95\\% confidence. Besides the change in 95\\% limits, non-zero BH natal kicks are not needed for progenitors of M$_{\\rm He} \\le 17$ M$_\\sun$ in order to explain the observed properties of Cygnus\\;X-1, but become necessary for M$_{\\rm He} > 17.5$ M$_\\sun$ (see Figure\\;10). Also, we note that a relatively small change on the range of V$\\rm_{pec,postSN}$ affects the derived constraint on V$_{\\rm k}$ qualitatively.  \\begin{figure} \\begin{center} \\plotone{fig10.eps} \\end{center} \\caption{The same plot of 2D joint V$_{\\rm k}$--M$_{\\rm He}$ confidence levels as Figure\\;8, calculated with V$\\rm_{pec,postSN} \\le 10$ km s$^{-1}$ instead of the original range derived Section\\;5.} \\end{figure}  \\subsection{Super-Synchronized Orbit}  After considering several previous measurements of the BH companion's surface rotation speed (V$_{\\rm rot} \\sin i$), \\cite{caballero...09} adopted V$_{\\rm rot} \\sin i = 95 \\pm{6}$ km s$^{-1}$. \\cite{orosz...11} found that the ratio of the BH companion's spinning frequency to the orbital frequency ($f_{\\Omega}$) was $1.400 \\pm{0.084}$, which was derived based on their results of the inclination angle $i = 27^\\circ.06 \\pm{0^\\circ.76}$ and the companion radius R$_2 = 16.5 \\pm{0.84}$. This indicates that the BH companion is super-synchronized. We note that with the analysis presented here, we find none of our winning binaries have super-synchronized BH companions at the current epoch. They are all sub-synchronized with $f_{\\Omega}$ reaching $\\sim 0.87$ at maximum.  In an effort to examine how our standard assumptions can be modified and investigate whether super-synchronism is at all allowed by the models as indicated by the observations, we make two modifications to our analysis. We first remove the assumption that the pre-SN orbit is pseudo-synchronized, and randomly distribute the pre-SN spin of the BH companion between zero and its breakup angular frequency $\\rm \\Omega_c$. Next, we reduce the secular changes of the orbital parameters due to the influence of tides by multiplying the right hand side of Equation (21) -- (23) by a constant $f\\rm_{tide}$.  \\begin{figure} \\begin{center} \\plotone{fig11.eps} \\end{center} \\caption{The probability distribution functions of the ratio ($f_{\\Omega}$) of the companion's spinning frequency to the orbital frequency at current epoch, which are calculated with models of $f\\rm_{tide}$ = 0.2 (solid/blue), 0.5 (dashed/green), and 1.0 (dotted/red). All three set of models have the pre-SN spin of the companion randomly distributed between zero and the breakup angular frequency $\\Omega_{\\rm c}$.} \\end{figure}  As shown in Figure\\;11, by allowing the pre-SN spin of companion to be greater than pseudo-synchronization and keeping the tidal strength unchanged (i.e.  $f\\rm_{tide} = 1.0$), the maximum $f_{\\Omega}$ of the winning binaries increases to $\\sim 1.2$. Although it is getting close, this value is still below the observationally inferred one. Together with a weakened tidal strength of $f\\rm_{tide} = 0.2$ and 0.5, we comfortably find winning binaries with $f_{\\Omega} = 1.4$. Furthermore, the minimum pre-SN surface rotation speed of the companion in those winning binaries are $\\sim 500$ and 700 km s$^{-1}$ for $f\\rm_{tide} = 0.2$ and 0.5, respectively. Given the uncertainties in the physics of fast rotating massive stars, it seems that super-synchronism in Cygnus\\;X-1 is allowed by our models presented, if the companion is spinning faster than orbital pseudo-synchronization right before the core collapse event and the tides exerted on the companion are weaker than the nominal theoretical values.  \\subsection{Constraints on the BH Kick Direction}  Since the measured eccentricity of Cygnus\\;X-1 is very low (see Table\\;1), and the orbit has not been circularized much since the formation of the BH (see Figure\\;6), e$\\rm_{postSN}$ has to be very small too. As discussed in Section\\;8, we indeed find that e$\\rm_{postSN}$ falls within a range of $0.015 - 0.022$. From the orbital dynamics at core collapse, the very low e$\\rm_{postSN}$ might shed light on the properties of the natal kick imparted to the BH.  During the core collapse event, e$\\rm_{postSN}$ is affected by the amount of mass loss, the direction and magnitude of the natal kick, as well as e$\\rm_{preSN}$. As mass loss tends to increase the eccentricity, a natal kick in the right direction and magnitude is needed in order to counteract the effect of mass loss and result in a very low e$\\rm_{postSN}$. It turns out that with the observationally inferred constraints on M$\\rm_{BH}$, M$_2$, V$\\rm_{pec,postSN}$, and the extra constraints mentioned in Section\\;6, the eccentricities right after symmetric explosions (i.e. no natal kicks) are typically $\\sim0.1$, and overwhelmingly $< 0.45$. Also, the difference between pre- and post-SN eccentricities is mostly $< 0.15$. These link to the requirement of the small V$\\rm_{pec,postSN}$. Larger amounts of mass loss relative to the total mass of the pre-SN system not only leads to higher e$\\rm_{postSN}$, but also larger V$\\rm_{pec,postSN}$, and hence they are not allowed. For the same reason, the required kick velocity for getting the very low post-SN eccentricity also needs to be small. In addition, we note that under special conditions, e$\\rm_{postSN}$ could be low even though e$\\rm_{preSN}$ is high. \\cite{hills83} showed that if the symmetric supernova explosion occurs at the proximity of apastron, e$\\rm_{postSN}$ could be $\\sim 0$ for a specific range of mass loss.  As the eccentricity induced by mass loss is low, natal kicks do not have to contribute dramatically to make e$\\rm_{postSN}$ fall within the observationally required range. Although natal kicks have to be constrained in some directions, that constraint is not super narrow. By extracting data from the winning binaries, we found that 70\\% have $\\cos(\\theta) < 0$; the distribution in the azimuthal angle $\\phi$ for the kicks of the winning binaries does not deviate much from the a priori flat distribution. As a result, the requirement of a very low e$\\rm_{postSN}$ does constrain the natal kick directions, but the constraint is not particularly strong."
    },
    "1107/1107.5581_arXiv.txt": {
        "abstract": "{ While the SUSY flavor, CP and gravitino problems seem to favor a very heavy spectrum of matter scalars, fine-tuning in the electroweak sector prefers low values of superpotential mass $\\mu$. In the limit of low $\\mu$, the two lightest neutralinos and light chargino are higgsino-like. The light charginos and neutralinos may have large production cross sections at LHC, but since they are nearly mass degenerate, there is only small energy release in three-body sparticle decays. Possible dilepton and trilepton signatures are difficult to observe after mild cuts due to the very soft $p_T$ spectrum of the final state isolated leptons. Thus, the higgsino-world scenario can easily elude standard SUSY searches at the LHC. It should motivate experimental searches to focus on dimuon and trimuon production at the very lowest $p_T(\\mu )$ values possible. If the neutralino relic abundance is enhanced via non-standard cosmological dark matter production, then there exist excellent prospects for direct or indirect detection of higgsino-like WIMPs. While the higgsino-world scenario may easily hide from LHC SUSY searches, a linear $e^+e^-$ collider or a muon collider operating in the $\\sqrt{s}\\sim 0.5-1$ TeV range would be able to easily access the chargino and neutralino pair production reactions. } ",
        "introduction": "\\label{sec:intro}  The well-known instability of the scalar sector of the Standard Model (SM) to quadratic divergences is elegantly solved by the introduction of supersymmetry (SUSY)\\cite{review}. In the case of the Minimal Supersymmetric Standard Model (MSSM) with soft SUSY breaking (SSB) terms, the divergences in the scalar sector are rendered to merely logarithmic. Interplay between the electroweak sector and the SUSY partners suggests the superpartner masses should exist at or around the TeV scale to avoid re-introduction of fine-tuning.  While the MSSM may be very appealing, it does suffer several pathologies. Unfettered soft SUSY breaking terms lead to large rates for flavor-changing neutral current processes and $CP$ violation\\cite{masiero}. Inclusion of grand unified theories with SUSY may lead to unacceptably high rates for proton decay\\cite{pdecay}. And in gravity-mediated SUSY breaking models (SUGRA), gravitino production followed by late-time gravitino decays in the early universe are in conflict with the successful picture of Big Bang nucleosynthesis (BBN) unless re-heat temperatures after inflation are limited to $T_R\\alt 10^5$ GeV\\cite{gravprob}. The latter bound is in conflict with appealing baryogenesis models such as thermal\\cite{leptog} (or non-thermal\\cite{ntlepto}) leptogenesis, which require $T_R\\agt 2\\times 10^9$ GeV ($10^6$ GeV).  A common solution to the above four problems is to push the SUSY matter scalars into the multi-TeV regime\\cite{dine}. The heavy scalars thus suppress loop-induced flavor and $CP$ violating processes, and suppress proton decay rates. If the multi-TeV scalars derive from a SUGRA model with a simple form for the K\\\"ahler potential, then the gravitino mass $m_{3/2}$ is also expected to exist in the multi-TeV range. By pushing $m_{3/2}$ into the 10-50 TeV range, the gravitino lifetime can be reduced to $\\tau_{3/2}\\alt 1$ second, so the gravitino decays shortly before BBN begins, this solving the gravitino problem\\cite{moroi}.  At first glance, multi-TeV gravitino and scalar masses seem in conflict with SUSY electroweak fine-tuning. The possible SUSY electroweak fine-tuning arises from minimization of the scalar potential after electroweak symmetry breaking. Here, the tree-level electroweak breaking conditions are familiarly written as\\cite{wss} \\bea B\\mu=\\frac{(m_{H_u}^2+m_{H_d}^2+2\\mu^2 )\\sin 2\\beta}{2}\\;, \\label{eq:mssmB}\\ \\ \\ {\\rm and}\\\\ \\mu^2 =\\frac{m_{H_d}^2-m_{H_u}^2\\tan^2\\beta}{(\\tan^2\\beta -1)} -\\frac{M_Z^2}{2}  , \\label{eq:mssmmu} \\eea where $B$ is the bilinear SSB term and $m_{H_u}^2$ and $m_{H_d}^2$ are the up and down Higgs SSB masses evaluated at the weak scale, $\\mu$ is the superpotential Higgs mass term and $\\tan\\beta$ is the ratio of Higgs field vevs: $\\tan\\beta ={v_u\\over v_d}$.  A measure of fine-tuning \\be \\Delta_i\\equiv\\left|\\frac{\\partial\\log M_Z^2}{\\partial\\log a_i}\\right| \\ee was advocated in Ref. \\cite{barbieri}. More sophisticated measures were advocated in Ref's \\cite{diego}, while in Ref. \\cite{ccn}, the $\\mu$ parameter itself is taken as a measure of fine-tuning: the latter paper requires $|\\mu |\\alt 1$ TeV to avoid too much fine-tuning. This measure motivates the well-known hyperbolic branch/ focus point (HB/FP) region of minimal supergravity (mSUGRA or CMSSM) as allowing for heavy scalars with low $\\mu$ value and low fine-tuning\\cite{ccn,feng}. A virtue of the HB/FP region is that multi-TeV scalars can co-exist with apparent low levels of electroweak fine-tuning.  In this paper, we will consider supersymmetric models with large, multi-TeV scalar masses, but with low, sub-TeV superpotential $\\mu$ term. We consider the case with intermediate range gaugino masses. This scenario, with \\be |\\mu |\\ll m_{gauginos}\\ll m_{scalars} , \\ee has been dubbed ``higgsino-world'' by Kane\\cite{kane}, and leads to a sparticle mass spectrum with a light higgsino-like chargino $\\tw_1$ and two light higgsino-like neutralinos $\\tz_1$ and $\\tz_2$. In models with gaugino mass unification at $M_{GUT}$, then the state $\\tz_3$ will be mainly bino-like, while $\\tz_4$ and $\\tw_2$ will be wino-like.  While the higgsino-world scenario seems highly appealing due to its ability to reconcile multi-TeV scalars and gravitinos with low electroweak fine-tuning, it has perhaps fallen out of favor for two reasons. First, higgsino-world leads to a very low thermal relic density of neutralinos, not at all in accord with measurements from WMAP and other experiments which require\\cite{wmap7} \\be \\Omega_{\\rm DM}h^2=0.1123\\pm 0.0035\\ \\ \\ {\\rm at\\ 68\\%\\ CL} . \\ee Second, higgsino-world scenarios are not easily realized in the paradigm mSUGRA/CMSSM framework, since elevating scalar masses into the multi-TeV region for a given value of GUT scale gaugino mass $m_{1/2}$ pushes one beyond the HB/FP region into a portion of parameter space where radiative EWSB is not realized under the assumption of universal scalar masses $m_0$ at $M_{GUT}$.  Pertaining to the dark matter issue, a number of recent works have emphasized that the standard picture of a thermal SUSY WIMP as dark matter is subject to very high fine-tuning\\cite{bbox}. Furthermore, non-standard cosmologies have many desirable features, and may even be favored by string theoretic constructions. For instance, Kane {\\it et al.} have shown\\cite{lightmod} that at least one moduli field in string theory should maintain a mass at or around the 10 TeV scale. Such moduli fields can be produced via coherent oscillations in the early universe, and decay into WIMPs, thereby augmenting the WIMP abundance\\cite{mr}, or they can decay into SM particles, thus generating entropy and diluting the WIMP abundance. Gelmini {\\it et al.} have shown in this case that SUSY models with {\\it any} value-- either too high or too low-- of thermal WIMP abundance may give rise to the measured CDM abundance via the enhancement or diminution due to scalar field (moduli) decays\\cite{gg}. In particular, for higgsino-world with too low a thermal WIMP abundance, the light higgsino abundance can be enhanced by moduli decays, leading to the correct abundance of higgsino-like WIMPs.  Alternatively, in SUSY models wherein the strong $CP$ problem is solved by the introduction of Peccei-Quinn-Weinberg-Wilczek\\cite{pqww} invisible axion\\cite{ksvz,dfsz}, one must introduce an axion supermultiplet, which contains an $R$-parity even spin-0 saxion $s(x)$, along with an $R$-parity odd spin-${1\\over 2}$ axino $\\ta (x)$, in addition to the light pseudoscalar axion field $a(x)$. In models such as these, with a TeV-scale axino and a higgsino-like neutralino as LSP, the $\\tz_1$ abundance can be augmented by axino production and subsequent re-annihilation at temperatures above BBN but below neutralino freeze-out\\cite{ckls}. Depending on the various PQMSSM model parameters, the CDM consists of an axion/neutralino admixture, where either the axion or the neutralino can dominate the abundance\\cite{blrs}.  A third modifiation of the thermal WIMP abundance may also occur: a model with a supposed underabundance of neutralino dark matter may enjoy enhancement of the relic DM abundance due to thermal gravitino production\\cite{gravprod} followed by cascade decays to the LSP state\\cite{grav,bdrs}.  Pertaining to the issue of higgsino-world being difficult to realize in the paradigm mSUGRA model, we note that it is easily realized in models with non-universal GUT scale Higgs masses (NUHM)\\cite{nuhm}. In fact, in GUT models such as $SO(10)$, the matter supermultiplets live in the 16-dimensional spinor representation, while Higgs superfields live in 10 or other dimensional multiplets. In such models, there is little reason to expect matter-Higgs SSB universality at $M_{GUT}$.  For the above reasons, we feel that it may be opportune to reconsider the higgsino-world scenario, and whether such a scenario would be visible to LHC SUSY searches. Toward this end, we discuss in Sec. \\ref{sec:pspace} the higgsino-world parameter space and expected mass spectra. In Sec. \\ref{sec:Oh2}, we present calculations of the standard thermal neutralino abundance in the higgsino-world scenario, and discuss its direct and indirect detection in the case where non-standard cosmological processes augment the relic higgsino abundance. In Sec. \\ref{sec:lhc}, we evaluate the dominant sparticle production cross sections for the lighter matter states at the LHC and calculate their branching fractions. While higgsino production cross sections occur at possibly observable levels, the compressed spectra lead to sparticle decays with very low energy release, and very soft detectable particles. To the best of our knowledge, higgsino-world SUSY can effectively elude standard SUSY searches for jets plus missing $E_T$ $(MET)$, and also for isolated multi-leptons$+MET$ at LHC7 (LHC at $\\sqrt{s}=7$ TeV) with $\\sim 10$ fb$^{-1}$ of integrated luminosity. In Sec. \\ref{sec:ilc}, we discuss higgsino-world signatures at a TeV scale lepton collider such as ILC or a muon collider (MC). In Sec. \\ref{sec:conclude}, we present our final discussion and conclusions. ",
        "conclusions": "\\label{sec:conclude}  The higgsino-world SUSY scenario with multi-TeV scalars, $\\mu\\alt 250$ GeV and intermediate scale gauginos is very appealing in that it can reconcile a decoupling solution to the SUSY flavor, CP, $p$-decay and gravitino problems with apparently low levels of naturalness or electroweak fine-tuning. The scenario is characterized by a mass hierarchy $|\\mu |\\ll m_{1/2}\\ll m_0$, where $m_0$ is the GUT scale mass of matter scalars. The HW scenario is most easily realized in models with non-universal Higgs masses, where the weak scale values of $\\mu$ and $m_A$ are taken as free parameters. In the HW scenario, the $\\tw_1$, $\\tz_1$ and $\\tz_2$ states are all light with mass $\\alt 250$ GeV, and dominantly higgsino-like. The remaining sparticles may well be heavy and inaccessible to LHC searches.  The standard thermal abundance of higgsino-like $\\tz_1$ particles is well below WMAP-measured values. However, in appealing cosmological scenarios such as those containing TeV-scale scalar fields such as moduli, or in scenarios with mixed axion-$\\tz_1$ cold dark matter, the neutralino abundance can be easily pushed up into the measured range. If this is so, then there are excellent prospects for direct or indirect detection of higgsino-like relic WIMPs, and we expect experiments such as Xenon-100 or Xenon-1-ton to fully explore this possibility. Alternatively, in DM models such as the mixed $a\\tz_1$ scenario\\cite{blrs}, it is also possible to tune PQ parameters such that the WIMP abundance remains tiny, while the bulk of CDM is comprised of axions. Thus, the HW scenario will not be completely excludable by direct or indirect WIMP search experiments if no signals for WIMPs are seen.  At the LHC, gluino and squark production may be suppressed by large values of $m_{\\tg}$ and especially $m_{\\tq}$. The $\\tw_1\\tz_1$, $\\tw_1\\tz_2$, $\\tz_1\\tz_2$ and $\\tw_1^+\\tw_1^-$ production reactions are then dominant, but are difficult to detect at LHC due to the small $\\tw_1 -\\tz_1$ and $\\tz_2-\\tz_1$ mass gaps, which lead to very soft visible particle production. The reaction $pp\\to\\tz_1\\tz_2$ may lead to tightly collimated OS/SF dilepton pairs, although calculations of signal and background after simple cuts indicate these occur at unobservable levels. Trileptons from $\\tw_1\\tz_2$ production are also difficult to see due to the soft spectrum of isolated leptons. Our studies should motivate our experimental colleagues to push for di- and tri-muon analyses at the very lowest levels of $p_T(\\mu )$ which are possible.  A linear $e^+e^-$ collider such as ILC  or a $\\mu^+\\mu^-$ collider operating with $\\sqrt{s}\\sim 0.5-1$ TeV should be able to make a thorough search for the HW scenario. If HW SUSY is discovered at ILC, then it should be possible to extract the gaugino/higgsino content of the $\\tw_1$, $\\tz_2$ and $\\tz_1$ states using various kinematic and angular distributions along with beam polarization. Thus, the HW scenario provides a concrete realization of a SUSY construct which may well remain hidden from LHC and dark matter searches, but which is fully testable at a TeV-scale lepton collider."
    },
    "1107/1107.5254_arXiv.txt": {
        "abstract": "We studied flux emergence events of sub-granular scale in a solar active region. New Solar Telescope (NST) of Big Bear Solar Observatory made it possible to clearly observe the photospheric signature of flux emergence with very high spatial (0\\arcsec.11 at 7057~\\AA) and temporal (15~s) resolution. From TiO observations with the pixel scale of 0\\arcsec.0375, we found several elongated granule-like features (GLFs) stretching from the penumbral filaments of a sunspot at a relatively high speed of over 4~km~s$^{-1}$. After a slender arched darkening appeared at a tip of a penumbral filament, a bright point (BP) developed and quickly moved away from the filament forming and stretching a GLF. The size of a GLF was approximately 0\\arcsec.5\\, wide and 3\\arcsec\\, long. The moving BP encountered nearby structures after several minutes of stretching, and a well-defined elongated shape of a GLF faded away. Magnetograms from SDO/HMI and NST/IRIM revealed that those GLFs are photospheric indicators of small-scale flux emergence, and their disappearance is related to magnetic cancellation. From two well-observed events, we describe detailed development of the sub-structures of GLFs, and different cancellation processes that each of the two GLFs underwent. ",
        "introduction": "Flux emergence is the main source for the formation of magnetic structures in the photosphere. Magnetic flux tube emerges from the convective zone through the photosphere by means of magnetic buoyancy due to the lower gas pressure inside the flux tube \\citep{Par55}. Instead of an abrupt emergence of a large bundle of flux tube, \\citet{Zwa87} suggested that a large-scale magnetic structure forms through the coalescence of numerous small-scale magnetic elements. \\citet{Str99} observed that small-scale magnetic flux subsequently emerges to form an active region, and 3D magnetohydrodynamics simulations by \\citet{Che08,Che10} showed that small-scale magnetic elements uplifted by buoyant force gradually coalesce into larger structures such as pores and sunspots. Recently, \\citet{Sch10,Sch10b} observed the formation of an elongated granule at the vicinity of a sunspot with both photometric and spectropolarimetric data and suggested that emerging loops on small scales may play an important role in a developing sunspot with a penumbra.  Indeed, to study small-scale magnetic elements emerging in an active region is important to understanding the ovarall magnetic structures of a sunspot. Especially, flux emergence in a moat region has become an important topic in the context of the origin of fine-structures such as penumbral filaments and moving magnetic features (MMFs). Such fine-structures are believed to be magnetically connected with the large-scale active region field in the form of a sea-serpent configuration \\citep[etc.]{Har73, Sch02, Wei04}, and hence observational properties of those structures, such as magnetic configuration, magnetic flux content, velocity, and the direction of movements \\citep{Lee92,Shi01,Yur01,Zha03,Hag05} have become critical constraints of sunspot models.  Although we have acquired substantial knowledge about small-scale emerging flux owing to recent improvements in polarimetric sensitivity \\citep{Lit96, Dep02, Lit08}, it seems that many small-scale magnetic structures are still not fully detected and/or resolved \\citep[see for example,][]{ste11}. On the other hand, the most up-to-date photospheric photometric data show fine-structures on sub-granular scale that seem to be associated with magnetic fields \\citep{Goo10b}. Studying such photospheric fine-scale structures could provide a tool for understanding the behavior of unresolved magnetic fields. Emerging flux will is associated with a photospheric counterpart, such as elongated granules when a loop top passes the photosphere \\citep{Bra64, Sch10}, and many studies investigated morphological properties of such photospheric features in detail. For instance, \\citet{Str99} studied the pattern of granular motion in an emerging active region utilizing both continuum intensity and circular polarization data, and noted several interesting phenomena, such as a transient line-center (or continuum) darkening associated with upflows later followed by bright grains. Although the emergence events they studied were relatively large ($\\geq1$~Mm), as compared to those addressed in recent spectropolarimetric works, this study provided insights on photospheric signatures associated with the emergence of the magnetic field. Additional attempts utilizing up-to-date data showing features on sub-granular scale, even without detecting magnetic field, are worth performing to understand the process of flux emergence on small spatial scales.  Studies focused on photospheric fine-scale counterparts of the emerging flux in the moat region are not numerous, so far. Better understanding of fine-scale features in the photosphere produced by emerging flux will also provide important information on the magnetic structure itself. Our aim is to fill this gap by studying photospheric sub-granular scale features that are believed to be signatures of flux emergence. Here we present recent observations of photospheric granular motion in an active region associated with small-scale emerging flux observed with both high spatial resolution and time cadence. The data have been acquired with a 1.6~m clear aperture New Solar Telescope \\citep[NST;][]{cao10a,Goo10} operating at the Big Bear Solar Observatory (BBSO). ",
        "conclusions": "We observed the formation of elongated GLFs just outside a pore and a sunspot in an active region NOAA 11109. The high spatial and temporal resolution data were taken by the NST at BBSO with the aid of an adaptive optics under very good seeing conditions. Magnetic and chromospheric features that are associated with GLFs were examined using NST/IRIM and H$\\alpha$ data along with {SDO}/HMI data. Among several events of GLFs that we found in this active region, two events showed similar signatures in TiO images, however they displayed greatly different behavior in the H$\\alpha$ line.  The most significant findings are : 1) both elongated GLFs were preceded by elongated darkening and showed a sharp, dark boundary at their moving front. The footpoint behind that boundary was brighter than its surroundings; 2) one of the two events was large enough (3\\arcsec) to show a fine structure consisting of thin, dark and bright threads, which to the best of our knowledge were not reported before; 3) the footpoints of GLFs that have a magnetic polarity  opposite to that of the nearby sunspot (pore), and are located farther from the sunspot (pore), moved away from the sunspot with a relatively high speed of 5.0~km~s$^{-1}$ and 4.3~km~s$^{-1}$, respectively.  Our conclusions are summarized as below. \\begin{enumerate} \\item   We observed the development of small-scale elongated GLFs on sub-granular scale less than 1\\arcsec\\ in extent at the tip of penumbral filaments. The footpoints of GLFs were co-spatial with newly emergent moving magnetic flux with the footpoint farther from the sunspot having magnetic polarity opposite to that of the sunspot. The observed properties of these photospheric fine-scale features were consistent with emerging $\\Omega$-shaped magnetic loops. \\item   In both observed events, after 1--2 Mm of displacement of the moving footpoints, the magnetic flux at moving footpoints cancelled out pre-existing opposite flux. A similar photospheric evolution of these events showed very different associated chromospheric dynamics. We speculate that the reason may lie in the different topology of the associated fields. \\end{enumerate}  Magnetic elements in the photosphere have mostly been studied using observations obtained with a G-band filter, since  magnetic elements in G-band images are seen as BPs with a contrast much higher than that of the surrounding granules \\citep{Ber95}. Many studies have reported that the BPs observed in G-band are indicative of small-scale flux emergence into the photosphere and also frequently associated with chromospheric events, such as brightenings in H$\\alpha$ or CaII images \\citep{Ber01}. The TiO $705$~nm line used to obtain our photospheric dataset is known to be temperature-sensitive, too, and as a good diagnostic line for imaging of the umbral regions \\citep{Ber03}. Recent observations employing a TiO broadband filter on the BBSO/NST reported detection of tiny BPs that might be related to small-scale magnetic fields \\citep{Goo10b, Abr10, And10}. The comparison between co-temporal and co-spatial G-band and TiO image revealed that TiO data show the same features as G-band \\citep{Abr10}. Given the size of magnetic elements on sub-granular scale, 0\\arcsec.5--1\\arcsec\\, in our case, and their weak magnetic flux, BPs' magnetic field is often not detected with the current state-of-the-art magnetographs, and thus photospheric intensity maps, such as those obtained with G-band or TiO-lines may be more suitable for studying the dynamics of tiny magnetic elements. In our case, the width and length of one of the GLFs was about 0.5\\arcsec\\, and 1\\arcsec\\,, respectively. The associated magnetic field was not detected in a {SDO}/HMI magnetogram, but the NST/IRIM data showed the Stokes-V signal even with the loss due to the cross-talk.  Observations of elongated GLFs as a manifestation of flux emergence in an active region were previously reported by \\citet{Str99} and \\citet{Sch10} who studied an emergence of an active region and an area immediately adjacent to the sunspot penumbra, respectively. \\citet{Sch10} measured the Stokes profiles along the elongated granule to show the existence of a horizontal field on the granule's spine and the vertical field of two footpoints, supporting the idea of the emergence of an $\\Omega$--shaped loop. Comparing to their works, our data show the entire evolution of a clear morphology of elongated GLFs with the unprecedented spatial and temporal resolution that has not been achieved from other observational data so far. Such observations on the morphology of the photospheric small-scale features could give us an important clue to understanding the interaction between emerging magnetic field and the granular convective flows.  Recent simulations on flux emergence show quite realistic photospheric phenomena, such as the formation of elongated granules and their associated transient darkening or BPs \\citep{Che08, Che10}. Our observed data show even more complex structures, such as dark and bright threads, a dark boundary at the moving front, and as well as a relatively high speed of migration of the footpoints of GLFs. Those complex structures of the emerging flux may indicate that even the small-scale flux is already fragmented. This finding was possible due to high spatiotemporal resolution. The time duration of the GLF's development was only $\\sim 20$ minutes, and the timespan during which we could see these fine threads was only $\\sim 5$ minutes. Therefore, low time-cadence may not catch such a development of fine structures. Also, the mixture of thin, dark and bright threads may be smeared out with the lower spatial resolution.  Observed fine-threads from the first GLF are similar to those of penumbral filaments. Moreover, both of two events occurred at tips of penumbral filaments and were associated with MMFs at relatively high speed compared to that of the moat flow. These findings lead us to fundamental question on the formation mechanism of GLFs and their relation with penumbral magnetic field. Although it is not certain whether the observed GLFs are direct extensions of attached penumbral filaments or not, the possible relation between developing GLFs and forming penumbral filaments is worth considering. In the moving flux tube model suggested by \\citet{Sch02}, magnetic field obtains the sea-serpent configuration as outward migrating waves are developed. \\citet{Wei04} suggested that the `turbulent pumping by granular convection' drags flux tube downward in the moat region to form sea-serpent configuration. Such a magnetic configuration is frequently adopted to interpret observation properties of MMFs and their relation with penumbral structures \\citep{Sai05, Sai08}. Our observed MMFs and GLFs could be also explained by the emergence of migrating $\\Omega$-shape crests of sea-serpent field and associated intensity enhancement due to radiative cooling.  Not only the formation of elongated GLFs themselves, but also the observation of two different cancellation phenomena merits discussion. Despite the similar photospheric observational aspects of formation of GLFs, the first one cancelled out without much of a noticeable chromospheric event, while the other one produced a significant chromospheric brightening and a jet. A brightening was also detected in {SDO}/AIA 1600~\\AA\\, images, while no signatures were observed in other wavelengths. This may indicate the magnetic reconnection probably involved small-scale field lines lying at a chromospheric height. Along with such brightening in both chromospheric H$\\alpha$ and EUV lines, the dark jet in H$\\alpha$ line is the typical indications of magnetic reconnection. Note that there is a negative pore in the upper-left corner of each panel in Figure~\\ref{mmf2_tio.irim.hal}, and the positive polarity of the GLF moves toward this pore. Dark fibrils lining up toward the lower-left from the pore seen in H$\\alpha$ centerline images indicate that the magnetic field of the negative pore connects to the positive polarity farther away from the area of interest. The direction of the dark jet that is seen in the H$\\alpha$ blue-wing images coincides with the field direction indicated by fibrils. This is consistent with a picture of magnetic reconnection occurring between a pre-existing large-scale field and a newly emerging $\\Omega$--shape loop of which one magnetic footpoint approaches the pre-existing opposite polarity. Such reconnection may produce both a dark jet with cool material ejecting upwards following the direction of the large-scale field lines and a brightening in either part of the approaching loops \\citep{Yur10}. The converging motion of the two opposing magnetic polarities is often believed to play an important role in driving magnetic reconnection \\citep{Chae02,von06}. The movement of the farther magnetic footpoint of observed GLFs was quite fast reaching $4$--$5$~km~s$^{-1}$. That rapid converging motion may have driven magnetic reconnection producing significant brightening in H$\\alpha$ in spite of the small size of magnetic flux. Such fast speed is not unexpected, and \\citet{Chae10} found a converging speed of H$\\alpha$ bright footpoints at 4.6~km~s$^{-1}$ before magnetic reconnection.  The absence of noticeable chromospheric activity associated with the first event suggests that it is possible that the disappearance of two closely located opposite polarity elements may not always be due to magnetic reconnection, but rather due to either emergence of U-loops or submergence of $\\Omega$-loops. The first event showed both the appearance of positive magnetic polarity (P1) and its disappearance with the pre-existing negative polarity (N2). Note that the protrusion of negative flux (N1) from the penumbra was also observed when P1 appeared. We carefully suggest that this phenomenon may represent emerging $\\Omega$-loops based on the magnetic configurations, such as the nearest footpoint having the same magnetic polarity as the pore. And also, this event was associated with the Doppler blue-shift taken by SDO/HMI with the value of about $-890$~m~s$^{-1}$ when average quiet region was set to zero. Magnetic reconnection could also occur without producing observable jet, and such possibility may not be fully ruled out. Still, it seems possible to suggest an additional scenario that could explain observed features. In case of the disappearance of P1, N2 was comparable to P1 both in size ($\\sim2\\arcsec$) and in magnetic flux ($\\sim 10^{18}$Mx). Such a small flux is likely to be a part of sea-serpent field near the surface forming the sequence of U- and $\\Omega$-shaped loops. Moreover, the H$\\alpha$ intensity at P1 showed its maximum before P1 encountered with N2. Based on that idea, we suggest an additional explanation for this event, such as the magnetic field of the newly emerging bipole and that of the pre-existing negative pole may be connected below the surface in the form of the sea-serpent field, and the emergence of such a field could result in the emerging bipole and the subsequent cancellation."
    },
    "1107/1107.2642_arXiv.txt": {
        "abstract": "{ We set up cosmological perturbation theory and study the cosmological implications of the so-called ``generalized Galileon'' developed in \\cite{Deffayet:2011gz,horndeski}. This is the most general scalar field theory whose Lagrangian contains derivatives up to second order while keeping second order equations of motion, and contains as sub-cases $k$-inflation, $G$-inflation and many other models. We calculate the power spectrum of the primordial curvature perturbation, finding a modification of the usual consistency relation of the tensor-to-scalar ratio in $k$-inflation or perfect fluid models. Finally we also calculate the bispectrum, which contains no new shapes beyond those of $k$-inflation. }   \\begin{document} ",
        "introduction": "Recently a great deal of effort has been devoted to developing consistent modified gravity theories which, as an alternative to to dark energy or the cosmological constant, may provide an explanation of cosmic acceleration (for a recent review, see \\cite{Clifton:2011jh}). Such IR modification of gravity arise, for example, in higher-dimensional setups as in the DGP model where self-acceleration is sourced by a scalar field $\\phi$, the helicity-0 mode of the 5D graviton. On small scales the DGP model also reproduces general relativity due to non-linear interactions operating through the Vainstein mechanism \\cite{Vainshtein:1972sx}. Other theories of modified gravity, such as DBI-galileons can also be obtained from a higher dimensional approach \\cite{brane_ga}.  Generally, modified gravity models contain additional degrees of freedom which can often be viewed as scalar fields. In this paper we consider models with a {\\it single} scalar field $\\phi$, though allowing for couplings between $\\phi$ and the metric which may be very different from the standard couplings and potentials considered in many inflationary models.   Indeed, the context of our work is that of the ``generalized Galileon'' derived in \\cite{Deffayet:2011gz} (see also \\cite{horndeski}).  While there are numerous different models of inflation with different actions and potentials \\cite{inf_model} -- for example chaotic inflation, small/large-field inflation, $k$-inflation and DBI-inflation to name a few -- single-field $k$-inflation \\cite{kess}, in particular, attracted a lot of attention since it is the most general scalar field theory with a Lagrangian containing derivatives up to \\emph{first} order $\\mathcal{L}=\\mathcal{L}(\\phi,\\nabla\\phi)$ (and hence, clearly, having second order equations of motion).  Indeed, all the inflationary models listed above are a special case of $k$-inflation, and for that reason the development of  cosmological perturbation theory in general case of $k$-inflation has a broad range of applicability.  In this paper we go beyond $k$-inflation and study the most general Lagrangian in 4 space-time dimensions which contain both first {\\it and second} derivatives of the scalar field, $\\mathcal{L}=\\mathcal{L}(\\phi,\\nabla\\phi,\\nabla\\nabla\\phi)$, but constructed  such a way that the field equations for both the metric and the scalar field remain second order \\cite{horndeski,Deffayet:2011gz}. This prevents the theory from having extra degrees of freedom as well as Ostrogradski instabilities \\cite{Ostrogradski},  and thus yields a possibly viable higher-order derivative scalar field theory. This general Lagrangian includes, in certain limits, the decoupling limit of DGP model \\cite{Dvali:2000hr} as well as some consistent theories of massive gravity \\cite{dgpmass,deRham:2011by}; and it also includes the ``Galileon'' model \\cite{Nicolis:2008in}, the conformal Galileon \\cite{conf_gal}, K-mouflage \\cite{Babichev:2009ee}, and also ``$G$-'' or ``KGB'' inflation \\cite{Deffayet:2010qz,g-inf}. Finally, in a very simple limit (which can be useful to check results) it also reduces to $k$-inflation. Our aim is to study cosmological perturbation theory in the context of this very general Lagrangian, outlined in Section \\ref{sec:model}, working both to second as well as third order so as to study non-Gaussianities. The formalism we develop is therefore applicable to a large class of models. As such, this paper is rather technical -- since we crank the standard handle of cosmological perturbation theory -- but the results are rather general.   As we will see below, see equation (\\ref{action}), the Lagrangian we consider \\cite{Deffayet:2011gz} contains {\\it four free functions} which will play an important r\\^ole in the following. We denote them by \\be K(X,\\phi),\\qquad \\text{and}\\qquad G^{(n)}(X,\\phi), \\qquad n=1,2,3 \\ee where $\\phi$ is the scalar field, and \\be X=-\\frac{1}{2}g^{\\mu \\nu}\\partial_\\mu \\, \\phi \\partial_\\nu \\phi, \\label{Xdef} \\ee is its kinetic term (we work with a mostly positive signature).  To obtain the Lagrangian of $k$-inflation \\cite{kess} one must simply set \\be K\\neq 0 \\; ,  \\qquad G^{(1,2,3)}(X,\\phi) = 0, \\qquad \\qquad (\\text{$k$-inflation limit}). \\label{kinflation-limit} \\ee The  Galileon Lagrangian \\cite{Nicolis:2008in,cdcov} is obtained when \\be K = X c_{(0)},\\qquad   G^{(1)} = Xc_{(1)},\\qquad G_{,X}^{(2,3)}=Xc_{(2,3)},  \\qquad \\qquad ({\\rm Galileon \\; limit}) \\label{Galileon-limit} \\ee where the $c_{(n)}$ are (dimensionful) constants. (The effective scalar field Lagrangian in the decoupling limit of the DGP model is obtained by setting $c_{(2)}=c_{(3)} = 0$.) In this case the Lagrangian will be invariant under shifts of the field, namely $\\phi\\rightarrow \\phi+c$. In general, when $K(X,\\phi)$ and $G^{(n)}(X,\\phi)$ are non-vanishing functions of $X$ and $\\phi$, shift symmetry is broken (as indeed is the ``Galileon'' symmetry \\cite{Nicolis:2008in} since we work in curved backgrounds). Without shift symmetry, questions of the stability of the theory against large renormalisation may arise \\cite{bdst_ng}, but we do not dwell on this point here.  To obtain the Lagrangian of $G$-inflation \\cite{g-inf} (or equivalently KGB-inflation \\cite{Deffayet:2010qz}), one must set \\be K,  G^{(1)} \\neq 0 , \\qquad G^{(2)}=  G^{(3)} = 0 \\; \\qquad \\qquad \\text{($G$-inflation limit}). \\label{Ginflation-limit} \\ee Cosmological perturbations has been studied in Galileon model \\cite{defelice,kobayashi,g-inf,lss_gal,galileon_inf,kyy2,Charmousis:2011bf,Gubitosi:2011sg}. Primordial non-Gaussianities \\cite{defe_ng,mk_ng,bdst_ng,cdmnt_ng,kyy1,renaux} have been studied in $G$-inflation limit, and it was shown \\cite{cdmnt_ng} that to leading order in slow-roll parameters there are no new momentum shapes for the bispectrum (beyond the standard ones of $k$-inflation). Despite that, a distinctive feature of $G$-inflation is that, though it has only a single scalar field, its energy momentum tensor takes the form of an imperfect  fluid \\cite{Deffayet:2010qz,Pujolas:2011he} (contrary to the $k$-essence models).  This fact sheds some light on the understanding of imperfect fluids in cosmology, and furthermore can lead to violations of the null enery condition thus opening up a whole range of possibly exotic phenomena \\cite{cnt}.  Problems of superluminality may arise in this case, though it is still under debate as to whether causality is violated in general \\cite{Babichev:2007dw}. Previous understanding of the conservation of curvature perturbation on super-Hubble scales relied highly on $k$-essence or perfect fluid models: it is interesting to see \\cite{Naruko:2011zk,Gao:2011mz} that this conservation law still holds in the generalized Galileon model, that is when the functions $K$ and $G^{(n)}$ are all taken to be arbitrary functions of $\\phi$ and $X$.    In this work, we keep the greatest generality by taking the full Lagrangian of generalized Galileon \\cite{Deffayet:2011gz,horndeski} into account: that is, arbitrary functions $K$ and $G^{(n)}$. We derive the background equations of motion for an inflationary background, around which we calculate the power spectrum and the bispectrum of the primordial curvature perturbation $\\zeta$ (the background evolution and linear perturbations of the generalized Galileon \\cite{Deffayet:2011gz,horndeski} were also investigated very recently in \\cite{kyy2}, see also \\cite{Charmousis:2011bf} for a discussion of self-tuning mechanism on FRW background). We find a modification of the familiar consistency relation of the tensor-to-scalar ratio. From the theoretical point of view, the results presented here can be used to study the evolution of cosmological perturbations as well as the observational implications numerous models.    This paper is organized as follows. In Sec.~\\ref{sec:model}, we review the generalized Galileon model. Background evolution of our model and linear perturbations are studied in Sec.~\\ref{sec:linear}. In Sec.~\\ref{sec:bispectrum}, we compute the full third-order action for the curvature perturbation and evaluate the corresponding contributions to the bispectrum. Final section is devoted to conclusion and discussion. ",
        "conclusions": "In this paper we have studied cosmological perturbation theory in the ``generalized Galileon\" model proposed in \\cite{Deffayet:2011gz, horndeski}. In 4 space-time dimensions this is the most general scalar field theory whose Lagrangian contains derivatives up to second order while keeping the equations of motion which are second order and lower. Our model, which includes $k$-inflation as a special case, represents a large class of scalar field models which has not been investigated so far.  The present work is the first step in exploring the cosmological implications of ``generalized Galileon\" models. By determining the most generic second and third order actions for the curvature perturbation, we calculated the power spectra of scalar and tensor perturbations as well as the bispectrum of the curvature perturbation. We found modifications of both the amplitude and propagation speed of tensor perturbations due to the presence of $\\mathcal{L}^{(2)}$ and $\\mathcal{L}^{(3)}$ Galileon terms. Correspondingly, the tensor-to-scalar ratio is modified relative to $k$-essence (see also \\cite{kyy2}). We have also showed that, although there are higher-order derivatives in the Galileon model relative to $k$-inflation, there are no new contributions to the bispectrum of the curvature perturbation. In order to get a feel for the strength of the bispectrum, we evaluated the non-linear parameter $f_{\\text{NL}}^{\\text{equil}}$ for the equilateral configuration $k_1=k_2=k_3$.  In the future, the strength and shape of primordial non-gaussianities will be constrained both by CMB data as well as, independently, by galaxy clustering data.  Indeed, it is expected  \\cite{LSS} that upcoming observations of high redshift clusters may possibly enable one to put limits on $f_{\\text{NL}}$ at the level of a few tens. It remains a very interesting and open question to see whether, as a result of this different data, strong constraints may be put on the different unknown functions appearing in the generalized Galileon model. Our result (\\ref{B_zeta}) will be a starting point for such an analysis."
    },
    "1107/1107.4064_arXiv.txt": {
        "abstract": "\\par We present 21-cm observations and models of the {\\sc H\\,i} kinematics and distribution of NGC 4244, a nearby edge-on Scd galaxy observed as part of the Westerbork Hydrogen Accretion in LOcal GAlaxieS (HALOGAS) survey. Our models give insight into the {\\sc H\\,i} kinematics and distribution with an emphasis on the potential existence of extra-planar gas as well as a negative gradient in rotational velocity with height above the plane of the disk (a lag).  Our models yield strong evidence against a significantly extended halo and instead favor a warp component along the line of sight as an explanation for most of the observed thickening of the disk.  Based on these models, we detect a lag of $-9^{+3}_{-2}$ km s$^{-1}$ kpc$^{-1}$ in the approaching half and $-9\\pm$2 km s$^{-1}$ kpc$^{-1}$ in the receding half. This lag decreases in magnitude to $-5\\pm$2 km s$^{-1}$ kpc$^{-1}$ and $-4\\pm$2 km s$^{-1}$ kpc$^{-1}$ near a radius of 10 kpc in the approaching and receding halves respectively.  Additionally, we detect several distinct morphological and kinematic features including a shell that is probably driven by star formation within the disk. ",
        "introduction": "\\label{Section1}   \\par Understanding the distribution and kinematics of vertically extended gas in disk galaxies is crucial for comprehending the relevance for the evolution of disk-halo flows (e.g.\\ the galactic fountain model presented by \\citealt{1976ApJ...205..762S} and refined by \\citealt{1980ApJ...236..577B}; and the chimney model of \\citealt{1989ApJ...345..372N}), interactions with the intergalactic medium (IGM) such as cold accretion of primordial gas (\\citealt{2005MNRAS.363....2K}, \\citealt{2008A&ARv..15..189S}) and interactions with neighboring galaxies.  Furthermore, comprehension of the interactions between ISM energization from star formation and accretion, i.e. feedback (e.g.\\ \\citealt{2000MNRAS.317..697E}) is also imperative.  \\par Galactic gaseous halos have been been observed in X-rays (e.g.\\ \\citealt{2006A&A...448...43T}), radio continuum (e.g.\\ \\citealt{1999AJ....117.2102I}), dust (e.g.\\ \\citealt{1999AJ....117.2077H}), ionized hydrogen (e.g.\\ \\citealt{2003A&A...406..493R}) as well as neutral hydrogen ({\\sc H\\,i}) (e.g.\\ \\citealt{1997ApJ...491..140S}, \\citealt{2007AJ....134.1019O}).   For the radio continuum (e.g.\\ \\citealt{2006A&A...457..121D}), diffuse ionized gas (DIG) (e.g.\\ \\citealt{2003A&A...406..493R}), X-ray (e.g.\\ \\citealt{2006A&A...457..779T}), and dust \\citep{1999AJ....117.2077H} there is a correlation between the presence of these halo components and star formation in the disk, both locally, in regions throughout the galaxy as well as globally.  This favors disk-halo flows akin to those in the aforementioned galactic fountain and chimney models as likely origins.  \\par Despite the likely connection between several halo components and star formation, the prevalence and properties of extra-planar {\\sc H\\,i} in galaxies and any potential connection to star formation is not well understood.  While there are many examples of {\\sc H\\,i} shells and holes associated with star formation (e.g.\\ \\citealt{2008A&A...490..555B}), no correlation between the prevalence of widespread {\\sc H\\,i} halos and star formation activity has been established among multiple galaxies as has been done for other halo components. Clearly, this can only be remedied through greatly expanded observations.  \\par Another clue to halo origins lies in their kinematics.  Trends concerning extra-planar {\\sc H\\,i}, such as the presence and magnitude of any decreases in rotation speed with height (lags) or correlations with the kinematics of DIG layers are useful in determining its origins.  One observational issue is whether {\\sc H\\,i} and DIG halos show the same kinematics, suggesting a common origin.    Recent measurements have shown lags in multiple DIG layers \\citep{2007ApJ...663..933H}, leading to various models which attempt to understand this gradient in terms of disk-halo flows and accretion of primordial gas.  \\par It has been shown that entirely ballistic models of disk-halo cycling, simulated in \\citet{2002ApJ...578...98C} and \\citet{2006MNRAS.366..449F}, are too simple in that they are unable to reproduce the observed kinematics.  It is possible the observed motions may partially be due to pressure gradients, magnetic tension \\citep{2002ASPC..276..201B} or external influences and feedback such as those described in \\citet{2008MNRAS.386..935F}.  Additionally, ``baroclinic\" (where gas pressure does not depend on density alone) hydrostatic models considered in \\citet{2006A&A...446...61B}, have been able to reproduce the observed lag in NGC 891, as has the hydrodynamic simulation of disk formation by \\citet{2006MNRAS.370.1612K}. It is likely that halos form and evolve via a combination of disk-halo flows as well as external influences.  \\par Of further interest is the radial variation of lags.  Basic geometric arguments predict a shallowing of the lag with increasing radius as may be seen in Figure 3 of \\citet{2002ApJ...578...98C}, but this could be complicated, for example, by pressure gradients \\citep{2002ASPC..276..201B}.  This will be discussed in greater detail in $\\S$~\\ref{Section6.2}.  \\par Detailed kinematic modeling of deep observations of {\\sc H\\,i} emission has been done for only a small number of galaxies including NGC 891 (\\citealt{1997ApJ...491..140S}, \\citealt{2007AJ....134.1019O}), NGC 5746 \\citep{2008ApJ...676..991R}, and NGC 2403 (\\citealt{2002AJ....123.3124F}). Through the Westerbork Hydrogen Accretion in LOcal GAlaxies (HALOGAS) survey \\citep{2011A&A...526A.118H},  which targets 22 edge-on or moderately inclined spiral galaxies for deep 10$\\times$12 hour observations using the Westerbork Synthesis Radio Telescope (WSRT), we aim to substantially increase this number in order to ascertain information concerning the origins of neutral extra-planar gas.  Through deep observations and kinematic modeling of these galaxies, we will establish whether there is any connection between extra-planar gas and star formation as well as the degree to which contributions from external origins are relevant.  \\par As part of the HALOGAS survey, we present observations and models of NGC 4244, a nearby edge-on Scd galaxy, which exhibits a thickened {\\sc H\\,i} layer upon inspection of the zeroth-moment map (presented here, Figures~\\ref{fig1},~\\ref{fig2}) as well as in previous work \\citep{1996AJ....112..457O}.  The first goal of the modeling is to determine the cause of the observed thickening, whether it is due to the presence of an {\\sc H\\,i} halo, flare, warp along the line of sight, or any combination of these possibilities. Second, we constrain whether a lag exists. Finally, we will identify local features, such as possible energy injection sites, which will be discussed further in $\\S$~\\ref{Section6.4}. This work will aid in establishing trends concerning the presence of halos and lags and any connection between disk-halo interactions and/or accretion.  In the case of NGC 4244, which has a low star formation rate (0.12 M$_{\\odot}$ yr$^{-1}$; \\citealt{2011A&A...526A.118H}), we are especially concerned with the potential existence of a lag.  If the trend found by \\citet{2007ApJ...663..933H} for DIG layers in galaxies with low star formation rate resulting in steeper lags holds true for {\\sc H\\,i}, then a steep lag is expected.  \\begin{figure*} \\epsscale{1} \\figurenum{1} \\includegraphics[width=170mm]{figure1.png} \\caption{The {\\sc H\\,i} zeroth-moment map overlaid on H$\\alpha$ (top) and MIPS 24 $\\mu$m (bottom) images.  Each image is rotated counter-clockwise by 43$\\,^{\\circ}$. {\\sc H\\,i} contours begin at $6.4\\times10^{19}\\mathrm{cm^{-2}}$ and increase by factors of 2.  The {\\sc H\\,i} beam is shown in the lower left-hand corner in white. The H$\\alpha$ image is supplied by Rene Walterbos from \\citet{1999ApJ...522..669H}. \\label{fig1}} \\end{figure*}  \\begin{figure} \\begin{center} \\figurenum{2} \\includegraphics[width=80mm]{figure2.png} \\caption{A zeroth-moment map of NGC 4244 showing the slice locations for the bv plots in Figure~\\ref{fig6} (solid), as well as the range of the evenly spaced slices in Figure~\\ref{fig10} (dashed).  Boxes a and b correspond to the regions shown in Figures~\\ref{fig11} and~\\ref{fig13} respectively. Contours begin at $6.4\\times10^{19}\\mathrm{cm^{-2}}$ and increase by factors of 2. The {\\sc H\\,i} beam is shown in the lower left-hand corner in white. An encircled dot and cross denote the approaching and receding halves respectively.\\label{fig2}} \\end{center} \\end{figure} ",
        "conclusions": "\\label{Section7}  \\par We present tilted ring models based on deep 21-cm observations of NGC 4244.  In these observations, we note a thickened {\\sc H\\,i} disk, with thinning at outer radii and evidence for a warp.  We also note asymmetries in the approaching and receding halves as well as localized features.  The models tested, in order of approximate increasing complexity, include a single disk with constant inclination, a single disk with a warp component along the line of sight, a single flaring disk, a thin and thick disk, and combinations of these.  \\par From these models, we conclude that the disk of NGC 4244 is substantially warped both perpendicular to as well as along the line of sight.  We do not detect an extended {\\sc H\\,i} halo, but with a scale height of 0.6 kpc, the disk is quite thick.  If {\\sc H\\,i} halos are due to disk-halo flows, then given the low star formation rate in NGC 4244 compared to other galaxies with {\\sc H\\,i} halos, the lack of a halo is unsurprising.  However, it is still to early to establish a trend, which will hopefully become apparent when additional cases are considered.  \\par In spite of not detecting an extended halo, we do detect a vertical lag in rotation speed of $-9^{+3}_{-2}$ km s$^{-1}$ kpc$^{-1}$ and $-9\\pm$2 km s$^{-1}$ kpc$^{-1}$ in the approaching and receding halves respectively. The magnitude of this lag decreases outward with radius to $-5\\pm 1$ km s$^{-1}$ kpc$^{-1}$ (approaching) and $-4\\pm 2$ km s$^{-1}$ kpc$^{-1}$ (receding) near 10 kpc, although subtle changes are within the error estimates.  Beyond 13 kpc, the S/N ratio no longer allows for a reliable lag estimate, so we cannot discern whether the lag vanishes entirely at some point.  \\par This is potentially the first instance where models favor a lag in {\\sc H\\,i} in a galaxy where an {\\sc H\\,i} halo does not exist.   There may be a lag in UGC 1281 \\citep{2011arXiv1103.0699K}, in which a halo is also absent, but in that case the lag cannot be distinguished from a warp component along the line of sight without deeper observations.  Given fluctuations in the position angle seen in the zeroth and first moment maps, it is possible that the warp in NGC 4244 begins closer to the center than in the models presented, which may indeed lessen or eliminate the need for a lag.  However, such a scenario is unlikely as these fluctuations may instead be due to localized features, a bar or spiral arms, which would be more in-line with current understanding of warps (\\citealt{1990ApJ...352...15B}, \\citealt{2011arXiv1101.1771V}).  \\par Modeling aside, we observe two substantial localized features which may correlate with star formation:  a shell-like feature and an unexplained curved structure originating in the midplane and possibly connected with a feature extending up to a height of 1.5 kpc.  A third distinct feature, a faint {\\sc H\\,i} streak is anomalous.  Finally, no additional substantial energetic features are detected, although numerous small-scale shell-like features are seen, but not displayed.  \\par We hope to use information gathered from our models of NGC 4244 in addition to information from the rest of the HALOGAS galaxies to further the development of trends concerning {\\sc H\\,i} kinematics, warps, extra-planar gas and any connections with disk-halo interactions, as well as accretion."
    },
    "1107/1107.3969_arXiv.txt": {
        "abstract": " ",
        "introduction": "} The field between the OB groups in Vela ($262^\\circ<l<268^\\circ$) and Car~OB1 ($284^\\circ<l<288^\\circ$) is not known to be dominated by any prominent OB association (see \\cite{humphreys78}, \\cite{melnik95}). The longitude range $283^\\circ$\\,-\\,$284^\\circ$ is thought to correspond to a tangential direction of a large segment of the Carina arm in both the 3- and 4-arm models of the grand design of the Milky Way \\cite{russeil03}. Our study is based on $uvby\\beta$ photometric distances and provides new insights on the apparent groupings and layers toward $281^\\circ<l<285^\\circ$ in the Galactic disk. This field clearly stands apart from the extended H{\\sc ii} features toward Car OB1 and contains several prominent H{\\sc ii} regions. ",
        "conclusions": ""
    },
    "1107/1107.3702_arXiv.txt": {
        "abstract": "{} {To investigate the dependence of the occurrence of active galactic nuclei (hereafter AGNs) on local galaxy density, we study the nuclear properties of $\\sim$ 5000 galaxies in the Coma Supercluster whose density spans two orders of magnitude from the sparse filaments to the cores of rich clusters. } { We obtained optical spectra of the nuclei of 283 galaxies using the 1.5m Cassini telescope of Bologna observatory. Among these galaxies, 177 belong to the Coma Supercluster and are added to the 4785 spectra available from SDSS (DR7) to fill-in the incomplete coverage by SDSS of luminous galaxies. We perform a spectral classification of the nuclei of galaxies in this region (with a completeness of 98\\% at $r\\leq$17.77), classifying the nuclear spectra in six classes: three of them (SEY, sAGN, LIN) refer to active galactic nuclei and the remaining three (HII, RET, PAS) refer to different stages of starburst activity. We perform this classification as recommended by Cid Fernandes and collaborators using the ratio of $\\lambda$ 6584 [NII] to $\\rm H\\alpha$ lines and the equivalent width (EW) of $\\rm H\\alpha$ (WHAN diagnostic), after correcting the last quantity by $1.3~\\AA$ for underlying absorption. } { We find that 482 (10\\%) of 5027 galaxies host an active galactic nucleus (AGN): their frequency strongly increases with increasing luminosity of the parent galaxies, such that 32\\% of galaxies with $\\rm Log (L_i/L_{\\odot})\\geq 10.2~ $ harbor an active galactic nucleus at their interior. In addition to their presence in luminous galaxies, AGNs are also found in red galaxies with $<g-i> \\simeq 1.15 \\pm 0.15$ mag. The majority of SEY and sAGN (strong AGNs) are associated with luminous late-type (or S0a) galaxies, while LIN (weak AGNs) and RET (\"retired\": nuclei that have experienced a starburst phase in the past and are now ionized by their hot evolved low-mass stars), are mostly found among E/S0as. The number density of AGNs, HII region-like, and retired galaxies is found to anti-correlate with the local density of galaxies, such that their frequency drops by a factor of two near the cluster cores, while the frequency of galaxies containing passive nuclei increases by the same amount towards the center of rich clusters. The dependence of AGN number density on the local galaxy density is greater than the one implied by morphology segregation alone. } {} ",
        "introduction": "It appears that 1.5m class telescopes, especially those ones located in zones of marginal climatic conditions, have lost much of their $raison~d'etre$, particularly after the advent of the Sloan Digital Sky Survey (SDSS). They however maintain a crucial didactic value for the training of young astronomers and provided that they are employed for suitable projects, can still make some significant niche contribution. One appropriate project in this respect is, in our opinion, the determination of the spectral properties of nuclei of bright galaxies, as part of a general study to perform a complete census of active galactic nuclei (AGNs). A spectral resolution on the order of R=1000 in the red channel (the blue-arm is more penalized by the relative blindness of CCDs) is sufficient to measure in a matter of a few minutes, even in non-photometric conditions, some crucial line ratios in the nuclear regions of galaxies, owing to their high surface brightness.\\\\ In this spirit, one of us (G.G.) invited yearly his first year master students to participate in some observing runs at the 1.5m Loiano telescope of the Observatory of Bologna (It). Unsurprisingly, many nights were affected by poor weather, although an amazing wealth of nuclear spectra (283) of nearby galaxies was accumulated, suitable for detecting any AGNs harbored in their interior.\\\\ A large amount of data (177) were gathered for the region of the Coma Supercluster, on which this work is focused. Following a fortunate run of 7/8 photometric nights in spring 2011 and the appearance of Mahajan et al. (2010), who investigated the active galactic nucleus (AGN) environmental dependence in the Coma Supercluster relying purely on SDSS data, regardless of its well-known incompleteness at the bright  end (Blanton et al. 2005a,b,c), we decided to publish our measurements derived in the past six years, with the hope that they might contribute to improving our understanding of AGNs.\\\\ There is a clear need to establish the frequency of AGNs of various types, in various environments, locally and in a cosmological perspective, to improve our understanding of galaxy assembly. Nevertheless, even after the advent of SDSS, which has permitted many extensive studies of galaxies, the processes involved remain unclear. Hao et al. (2005) estimated that 4-10\\% of all galaxies in the SDSS harbor an AGN; Kauffmann et al. (2003) estimated that up to 80\\% of galaxies more massive than $10^{11} M_\\odot$ host a supermassive black hole, either dormant or active. The role of the environment in triggering or inhibiting nuclear activity remains disputed. Kauffmann et al. (2004), analyzing $\\sim 122000$ SDSS spectra, found a decrease in the fraction of powerful AGNs (with high [OIII] luminosity, i.e. dominated by Seyferts) with increasing galaxy environment density, mainly because the most powerful AGNs are hosted by late-type galaxies, which are the least likely to exist in dense environments. Popesso \\& Biviano (2006) found that the AGN fraction decreases with the increasing velocity dispersion of galaxies in groups and clusters, being higher in dense, low-dispersion groups, in contrast to the findings of Shen et al. (2007) and Miller et al. (2003) who found that the frequency of low activity AGNs does not correlate with environment. \\\\ The classification of AGNs based on optical spectra has been traditionally performed using the BPT diagnostic diagram of Baldwin et al. (1981), which adopts the measurement of at least four spectral lines: $\\rm H\\beta$, O[III], $\\rm H\\alpha$, and N[II]. The ratio $\\rm H\\alpha$/N[II] (where $\\rm H\\alpha$ must be corrected for any underlying stellar absorption, as stressed by Ho et al. 1997) differentiates AGNs from HII region-like nuclei, while the ratio $\\rm H\\beta$/O[III] allows us to separate the strong AGNs from the weaker LINERs. More recently Cid Fernandes et al. (2010, 2011) introduced a two-line diagnostic diagram, named WHAN, based on the $\\rm H\\alpha$/N[II] ratio combined with the strength of the $\\rm H\\alpha$ line (corrected for underlying stellar absorption) to discriminate both strong and weak AGNs, which are both supposed to be ionized by active black holes, from \"fake AGNs\", dubbed \"retired galaxies\", whose ionization mechanism is provided by their old stellar population. As argued by several authors (e.g. Trinchieri \\& di Serego Alighieri 1991, Binette et al. 1994, Macchetto et al. 1996, Stasi{\\'n}ska et al. 2008, Sarzi et al. 2010, Capetti \\& Baldi 2011), hot evolved stars, such as post-asymptotic giant branch stars and white dwarfs, can produce a substantial diffuse field of ionizing photons: it has been shown that they can produce emission lines with ratios mimicking those of AGNs and - in particular - of LINERs.  \\\\ An equivalent width (hereafter EW) of $\\rm H\\alpha$ of at least 3 \\AA ~is ultimately required for a nucleus to be considered ionized by a central black hole (Cid Fernandes et al. 2010, 2011). This quantitative threshold of course might be questioned. Among the galaxies in our sample, we might consider for example the elliptical NGC3862 (CGCG 97-127), which is the second brightest galaxy of A1367 and was classified as an AGN by V{\\'e}ron-Cetty \\& V{\\'e}ron (2006). It harbors the radio galaxy 3C-264 (Gavazzi et al. 1981) and an optical jet (Crane et al. 1993), leaving little doubt that it contains a central monster. Its optical spectrum (taken either at Loiano or from SDSS) displays a strong NII emission line (of EW 4.81 \\AA) and a weak $\\rm H\\alpha$ EW of 1.51 \\AA. After adding 1.3 \\AA~ to the underlying absorption, the galaxy just misses the  $\\rm H\\alpha>3~\\AA$ threshold, thereby fails to be classified as a LIN but not as a RET according to the criteria of Cid Fernandes et al. (2010, 2011)\\footnote{ Similar examples (see Section \\ref{results}) convinced us to lower the threshold for LIN from 3 \\AA~ to 1.5 \\AA~ in the present analysis.}.\\\\ The Coma Supercluster, on which this paper is focused, is an ideal laboratory for initiating a complete census of AGNs at $z=0$ because problems related to the incompleteness of SDSS are less severe here than elsewhere. Even its largest sized galaxies have diameters $<2-3$ arcmin, making the \"shredding\" problem much less severe than for other nearby environments (e.g. Virgo). The brightest galaxies have been well-studied and even the faintest galaxies belonging to the Coma Supercluster have sizes in excess of 10 arcsec, allowing their classification by visual inspection on SDSS plates. Moreover, at the distance of 100 Mpc in the Coma Supercluster, the three arcsec fibers adopted by SDSS provide truly \"nuclear\" spectra. Studying the Coma Supercluster using both the integrated and nuclear properties of galaxies at $z=0$ derived from SDSS data, allows us to compare galaxies properties for a variety of environmental conditions, from the center of two rich clusters to the sparse filamentary regions where the local galaxy density is $\\sim$ 100 times lower, without the biases caused by studying galaxies at different distances. It provides a complementary view to the many statistical analyses that have been made using SDSS data (e.g. Kauffmann et al. 2004). Combined with studies of galaxy evolution at higher redshift it will shed light on the mechanisms and processes that contribute to the evolution of galaxies in the cosmological context. \\\\ Owing to the two-line WHAN diagnostics, the red-channel spectra taken at Loiano can contribute significantly to the census of AGNs by filling the residual incompleteness of the SDSS spectral database for luminous galaxies caused by shredding and fiber conflict (Blanton et al. 2005a,b,c).\\\\ The layout of this paper is as follows. In Sect. \\ref{The data}, we illustrate the observations taken at the Loiano observatory and the data reduction procedures. The nuclear spectra are given and classified in Sect. \\ref{results}. The census of AGNs of the various types in the Coma Supercluster is carried out in Sect. \\ref{AGN COMA}. Our results are compared with Mahajan et al. (2010) and discussed in Sect. \\ref{Discussion}. The standard cosmology is assumed, with $Ho$=73 $\\rm km~sec^{-1}~Mpc^{-1}$. ",
        "conclusions": "\\label{Discussion} \\begin{figure} \\begin{center} \\includegraphics[width=9cm]{gavazzi_fig9.ps} \\caption{The frequency distribution of SEY ({\\bf top left}); sAGN ({\\bf top right}); LIN ({\\bf middle left}), HII ({\\bf middle right}); RET ({\\bf bottom left}); PAS ({\\bf bottom right}) and all AGNs ({\\bf top}) with RET (pink) and without RET (black) in 4 bins of local galaxy density. } \\label{dens} \\end{center} \\end{figure} \\begin{figure}% \\begin{center} \\includegraphics[width=9cm]{gavazzi_fig10.ps} \\caption{The frequency distribution of the ratio of SEY+sAGN to giant spirals (SOa-Sb); of LIN to giant galaxies (E-Sb); of RET to early-type galaxies (dE-S0a); of HII to SOa+late-type galaxies; and of PAS to early-type (dE-S0) galaxies in four bins of local galaxy density. } \\label{fracdens} \\end{center} \\end{figure} A dependence of AGN activity on local galaxy density similar to the one found in this work was proposed by Mahajan et al. (2010), who analyzed almost exactly the same sky volume. This apparent agreement is however the consequence of a manifold blunder in their analysis. First of all, the galaxy sample in the Coma Supercluster used by them suffers from 20\\% incompleteness at high luminosity (see Section 2) owing to the incompleteness of the SDSS spectral database, which is entirely disregarded in the Mahajan et al. (2010) analysis.\\\\ Secondly the  $\\rm H\\alpha$ spectral line measurements from the SDSS database have not been corrected for underlying stellar absorption. Thirdly the fraction of AGNs is computed by Mahajan et al. (2010) by normalizing to the total number of galaxies, irrespective of their morphology, Owing to the first omission and the adoption of the BPT diagnostic, Mahajan et al. (2010) failed to classify as AGNs an enormous number of objects with [NII] in emission and $\\rm H\\alpha$ in weak absorption, particularly ones present in the core of rich clusters, that would increase the frequency of AGNs in the densest environments. At the same time, they overestimated the number of AGNs among galaxies with [NII] and $\\rm H\\alpha$ in emission (that become HII like nuclei when corrected). These are predominantly found in the field. If Mahajan et al. (2010) had adopted the correct procedure in producing their BPT diagram, they would have concluded that the fraction of AGNs would increase, not decrease, with galaxy density. The reason why AGNs are missing in clusters is that most galaxies that are bona-fide AGNs according to the BPT diagnostic diagram become fake-AGNs (RET) according to the WHAN diagnostic diagram, making the whole issue of the influence of the environment on the formation of AGNs very slippery.\\\\ Using the AGN classification scheme of Kauffmann et al. (2003) based on their revised BPT diagram, Kauffmann et al. (2004) analyzed the properties of powerful AGNs among $\\sim 122000$ galaxies with $14.5<r<17.77$ and median redshift $z \\sim 0.1$, from the SDSS DR1. The subtraction of the underlying stellar continuum was performed by fitting the emission-line-free regions of the spectrum with a model galaxy spectrum. They concluded that powerful AGNs occur predominantly in massive galaxies with $>10^{10} ~M_\\odot$, another manifestation of \"downsizing\" (Cowie et al. 1996, Gavazzi et al. 1996, Fontanot et al. 2009). Powerful AGNs are found among late-type, blue, star-forming galaxies. The decrease in their frequency with increasing galaxy density is mostly due to the decrease in the frequency of galaxies that are able to host AGNs (morphology segregation), but - to a lesser extent - is due to a genuine extra decrease in AGN frequency in high density environments. Von der Linden et al. (2010) argued that galaxies in dense environments are less likely to host a powerful optical AGN or a star-forming nuclear region. The frequency of optical AGNs in early-type galaxies declines in high density environments (by approximately a factor of two between the field and the center of clusters). They interpreted this finding by arguing that AGNs are fueled by mass loss from evolved stars. The mean age of the red galaxies increases towards the cluster center and provides less fuel for the central black hole.\\\\ Our analysis, which adopts the criteria of Cid Fernandes et al. (2010, 2011) for the classification of nuclei, is in full agreement with both Kauffmann et al. (2004) and Von der Linden et al. (2010). We conclude that, even after eliminating the effects of morphology segregation (by counting the frequency of nuclear activity in the proper bins of Hubble type), there is a residual decline in the frequency of nuclear activity (either triggered by black holes or sustained by the star formation) with increasing galaxy density (see Fig. \\ref{fracdens}). At any luminosity, this decline is quantified in a factor of $\\sim$ 2 per interval of density contrast of approximately 20, going from the less dense filaments that form the local cosmic web to the densest cores of rich clusters. Whichever mechanism is responsible for the suppression of over star formation in galaxies, and the migration of galaxy morphologies toward earlier Hubble types in rich and dense environments (see Boselli \\& Gavazzi 2006, Blanton \\& Moustakas 2009, Fumagalli \\& Gavazzi 2008), it has far more severe consequences for their nuclear regions. We are not yet in the position to distinguish these two effects. However, we are in the process of obtaining a complete set of $\\rm H\\alpha$ imaging follow-up observations of the Coma Supercluster that is appropriate for distinguishing the nuclear from the extended star-formation in galaxies in this representative part of the local Universe.\\\\"
    },
    "1107/1107.5725_arXiv.txt": {
        "abstract": "Detection of gravitational waves from the inspiral phase of binary neutron star coalescence will allow us to measure the effects of the tidal coupling in such systems. These effects will be measurable using 3$^{\\text{rd}}$ generation gravitational wave detectors, e.g. the Einstein Telescope, which will be capable of detecting inspiralling binary neutron star systems out to redshift $z\\approx 4$.  Tidal effects provide additional contributions to the phase evolution of the gravitational wave signal that break a degeneracy between the system's mass parameters and redshift and thereby allow the simultaneous measurement of both the effective distance and the redshift for individual sources. Using the population of $\\mathcal{O}(10^{3}\\text{--}10^{7})$ detectable binary neutron star systems predicted for the Einstein Telescope the luminosity distance--redshift relation can be probed independently of the cosmological distance ladder and independently of electromagnetic observations. We present the results of a Fisher information analysis applied to waveforms assuming a subset of possible neutron star equations of state. We conclude that for our range of representative neutron star equations of state the redshift of such systems can be determined to an accuracy of $8\\text{--}40\\%$ for $z<1$ and $9\\text{--}65\\%$ for $1<z<4$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5808_arXiv.txt": {
        "abstract": "In the excursion set approach to structure formation initially spherical regions of the linear density field collapse if the average density contrast within them exceeds some critical value, $\\delta_c$. This model allows one to make predictions for several fundamental aspects of structure formation such as the mass function of dark matter halos, their merger histories and clustering properties. The value of $\\delta_c$ is often calculated from the spherical or ellipsoidal collapse model, which provide well-defined predictions given auxiliary properties of the linear tidal field at a given point. We use two cosmological simulations of structure growth in a LCDM cosmology to test a key assumption used in calculating $\\delta_c$: that the {\\em shapes} of the initial Lagrangian patches that eventually collapse (or proto-haloes) are spherical. Our results indicate that the vast majority of dark matter proto-haloes are non-spherical, and have minor-to-major axis ratios that vary from $\\langle a_3/a_1\\rangle \\sim 0.4$ at the galaxy mass scale to $\\sim 0.65$ for rich galaxy clusters. We show that this non-sphericity likely originates from the asymmetry of the linear tidal field which pushes material onto, or away from, local density maxima in the linear density field. We study the implications of these results for the collapse barriers for CDM halo formation inferred from the classic ellipsoidal collapse model. Our results indicate that the ``standard'' ellipsoidal collapse model commonly adopted in the literature does not provide a full account of the possible collapse thresholds for halo formation, since the model predictions depend sensitively on the assumed {\\em shape} of the primordial perturbation. We show that an improved model, which accounts for the intrinsic shapes of proto-haloes, provides a much more accurate description of their {\\em measured} minimum overdensities. ",
        "introduction": "\\label{sec:intro} \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: aludlow@astro.uni-bonn.de}  In standard models for the growth of large scale structure, such as the $\\Lambda$-cold dark matter (CDM) model, dark matter haloes form hierarchically from low-amplitude Gaussian density fluctuations seeded during inflation. Galaxy formation is assumed to take place as gas cools and condenses in the centers of these haloes which subsequently ``light-up'' the underlying large scale matter distribution \\citep[e.g.,][]{White1978}. Although the hydrodynamical processes inherent to the formation of galaxies are difficult to model reliably, they are largely irrelevant to the over-all formation and clustering of the dark matter haloes. A first step towards understanding the statistics of the galaxy distribution is thus a detailed understanding of the spatial distribution and clustering of their haloes.  Since CDM haloes are highly non-linear objects much of our understanding of their properties has come from direct N-body simulations of their formation. Although useful, these simulations are often difficult to interpret and our understanding of the structure formation process can be significantly enhanced by augmenting them with simple physical models for the collapse and virialization of dark matter haloes.  One successful analytic approach to modeling structure formation has become known as the Extended Press-Schechter formalism \\citep[see, e.g.,][]{PS1974,Bond1991}. Here one identifies collapsed objects at a given time with regions of the linear density field {\\em expected} to collapse based on some dynamical model, such as the spherical \\citep{Gunn1972} or ellipsoidal collapse model \\cite[e.g.,][]{Zeldovich1970,Icke1973,WhiteSilk1979,Hoffman1986,Lemson1993, Bertschinger1994,Bond1996}. Although \\citet{PS1974} originally used this model to calculate the abundance of collapsed objects as a function of mass and redshift, it can be extended to compute a variety of hierarchical clustering statistics, such as halo merger histories \\citep[e.g.,][]{Lacey1993, ShethLemson1999a,vandenBosch2002}, and spatial correlations \\citep{MoWhite1996,Mo1997,Catelan1998, ShethLemson1999b}. A related approach, the peaks formalism, relates the abundance of bound objects to the abundance of local density maxima above some threshold in the smoothed linear density field \\citep{Peacock1985,Bardeen1986,LudlowPorciani2010}.  A fundamental aspect of both the Extended Press-Schechter approach and the peaks formalism is the need to identify which regions of the linear density field are expected to collapse to form bound objects. Here one typically gains insight from dynamical models that can provide physically-motivated thresholds for the linearly extrapolated density that is required for collapse to occur.  One common threshold, based on the spherical collapse model \\citep{Gunn1972}, assumes that haloes can be identified in the smoothed linear density field by means of a single free parameter: the initial density contrast, $\\delta_{sc}$. The basic properties of the halo distribution mentioned above can then be deduced from the trajectories of the linear density field as the spatial scale of the filtering is varied. In this case the threshold for collapse at redshift $z_c$ (in an Einstein-de Sitter universe) is $\\delta_{sc}\\approx 1.686 \\ (1+z_c)$. This result is approximately valid for a large range of cosmologies and is independent of mass scale and initial density contrast. \\citet{Eke1996} provide useful analytic formulae for calculating the value of $\\delta_{sc}$ for cosmologies with $\\Omega_M+\\Omega_{\\Lambda}=1$.  Although the simple spherical model works well in a statistical sense its validity is questionable; it is well-known, for example, that perturbations in Gaussian random fields are inherently triaxial \\citep{Doroshkevich1970,Bardeen1986}. This led to modifications of the spherical model to in incorporate the dynamics of ellipsoidal collapse.  Building upon the classic work by \\citet{LyndenBell1964} and \\citet{LinMestelShu1965}, a number of authors have studied the effect of internal shear on the collapse of a homogeneous ellipsoid in a uniformly expanding background \\citep[e.g.,][]{Icke1973,WhiteSilk1979,Peebles1980,WatanabeInagaki1991,Lemson1993}. \\citet{EisensteinLoeb1995}  and \\citet{Bond1996} generalized this model to follow the collapse of an initially spherical perturbation in the presence of external shear. In this case, the sphere is distorted into an ellipsoid whose principal axes are parallel to those of the external tidal field. One can then ascertain the impact of {\\em external} tidal shear on the properties of collapsed regions. For example, \\citet{Monaco1997b} and \\citet{Sheth2001} have studied how ellipsoidal collapse modifies the mass function of dark matter haloes in the excursion set formalism \\citep[see also][]{LeeShandarin1998}.  In the ellipsoidal collapse model of \\citet{Bond1996} the collapse time depends explicitly on the local ellipticity, $e$, and prolaticity, $p$, of the tidal field. Less spherical regions of the linear tidal field must overcome additional tidal stretching and thus require higher overdensities to collapse. \\citet{Sheth2001} showed that the dependence of the collapse threshold on $e$ and $p$ could be well-approximated by the solution to \\begin{equation} \\frac{\\delta_{ec}}{\\delta_{sc}}=1+\\beta\\biggr[ 5(e^2\\pm p^2) \\biggr(\\frac{\\delta_{ec}^2}{\\delta_{sc}^2}\\biggr) \\biggl]^{\\gamma}. \\label{eq:ecthreshold} \\end{equation} Here $\\beta$ and $\\gamma$ are numerical parameters determined from fits to the model results. \\citet{Sheth2001} used the results of \\citet{Doroshkevich1970} to compute the most probable values of $e$ and $p$ as a function of smoothing scale. At any ellipticity, the most probable prolaticity is $p_{{\\rm mp}}=0$, and for $p=0$ they find $e_{{\\rm mp}}=\\sigma/\\delta/\\sqrt{5}$. Using these results they showed that \\begin{equation} \\delta_{ec}=\\delta_{sc}\\biggr[ 1 + \\beta\\biggr(\\frac{\\sigma(M)^2}{\\delta_{sc}^2}\\biggl)^\\gamma\\biggl], \\label{eq:ecbarrier} \\end{equation} with $\\beta=0.47$ and $\\gamma=0.615$. Using their assumptions, eq.~\\ref{eq:ecbarrier} can be expressed as a function of the tidal ellipticity by substituting $\\sigma(M)=(\\delta/\\sqrt{5}) \\ e_{{\\rm mp}}$. (Various model-predicted collapse boundaries have been compared directly with N-body simulations by \\citet{Robertson2009}.)  Adding to this, the mass distribution in the vicinity of the linear density peaks from which haloes form are inherently triaxial \\citep{Bardeen1986}. This result also applies to the Lagrangian regions which collapse onto those peaks to form haloes \\citep[e.g.,][]{Porciani2002b,LudlowPorciani2010}. This implies that, in the absence of any external shear, the {\\em internal} shear associated with the initially non-spherical shape of a collapsing region may impact its collapse time, and consequently, the average density required for complete collapse. These subtleties are often over-looked in simple ellipsoidal collapse models, such as those of \\citet{Bond1991} and \\citet{Sheth2001}.  The purpose of the present paper is to investigate how simple modifications to the classic ellipsoidal collapse model affect the predicted shape of collapse barriers. Our analysis focuses on two cosmological simulations of structure formation, the results of which we compare directly with predictions from the model. The rest of the paper is organized as follows. In Section~\\ref{sec:numerics} we introduce our simulations and briefly describe our main analysis techniques. Section~\\ref{sec:protohaloes} describes some of the basic characteristics of proto-haloes in our simulations, focusing on their inertial and tidal tensors. We use these results in Section~\\ref{sec:collapse} to motivate a closer look at the ellipsoidal collapse model, and in particular, how the predicted collapse barriers depend on the assumed shape of the collapsing region. The predictions are compared directly with the results of two cosmological simulations in Section~\\ref{sec:boundary}. Finally, we end with a brief discussion in Section~\\ref{sec:discussion} before summarizing our main results in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  In the ellipsoidal model the collapse threshold depends explicitly on the shape of the tidal field at a given point ${\\mathbf{x}}$; the inferred collapse threshold can then be compared with the {\\em true} overdensity at that location to decide whether or not collapse should take place. The approach of \\citet{Sheth2001} was to adopt, for every spatial location, the most probabilistic values of $e$ and $p$ in Gaussian random fields, resulting in a barrier height that depends explicitly on filtering scale: $e_{\\rm{mp}}=\\sigma({\\rm M})/\\delta/\\sqrt{5}$ (since $p_{\\rm mp}=0$ for any $\\sigma({\\rm M})$). The resulting barrier height depends on only one variable: the characteristic mass scale under consideration.  We have shown that relaxing the assumption that the assumption of sphericity for the collapsing Lagrangian regions leads to very different predictions for the collapse barriers at any $e$, $p$, or mass scale. But what is the characteristic shape of a proto-halo and how is it related to the external tides, $(e,p)$, the characteristic mass scale, the filtering choice, etc? In this section we describe two simple methods for estimating in advance the potential shapes of collapsing proto-haloes at any position ${\\mathbf{x}}$.  One possibility is to estimate the characteristic shape of collapsing proto-haloes directly from the mass dependence of proto-halo shapes shown in Figure~\\ref{Fig:initial_shapes_MassDep}. The best-fit relations given in eq.~\\ref{eq:balinear} can then be used to determine the {\\em average} proto-halo shape expected for a given filter volume (which corresponds to a well-defined halo mass scale). At any given smoothing scale, $\\sigma(M)$, there is then a well-defined boundary $\\delta_B(e,p,M)$ that can be used to select possible sites for collapse based on the properties of the tidal field at $\\mathbf{x}$.  In general, this requires a calculation of the full ellipsoidal collapse barrier at every characteristic mass scale, since the axis ratios are themselves smooth functions of mass. In order to make the calculation simpler we have assumed that $p=0$ for all $\\mathbf{x}$ (so that the tidal field is fully characterized by its ellipticity, $e$) and that dark matter proto-haloes can be approximated as ``perfectly triaxial ellipsoids'', i.e. $a_2/a_1=a_3/a_2$, so that a single axis ratio is sufficient for describing their shape.  In Figure~\\ref{Fig:FracAboveMass} we plot the fraction of proto-haloes whose linear density contrast exceeds the threshold for collapse based on the method outlined above. Here we estimate, at any given halo mass scale, the typical value of $a_2/a_1$ from eq.~\\ref{eq:balinear}, set $a_3/a_1=(a_2/a_1)^2$, and compute the corresponding barrier as a function of tidal ellipticity in the usual fashion, by assuming $p=0$. At any given mass scale, the fraction of points that lie above this threshold are shown as solid circles connected by lines. Triangles connected by dashed lines show the fraction of haloes that lie above the barrier advocated by \\citet{Sheth2001}. We use blue and red lines to distinguish results obtained from our 150\\Mpch{} and 1200\\Mpch{} box runs, respectively.  Taking into account the average mass-dependence of proto-halo shapes we find that $\\simgt 90\\%$ now lie above the ellipsoidal collapse boundary. Adopting the \\citet{Sheth2001} barrier instead yields very different results: at $\\sim 10^{11} h^{-1} {\\rm M}_{\\odot}$, for example, only about 70\\% of proto-haloes should collapse given their averaged ellipticities and density contrasts. At the highest mass scales, where haloes are initially more spherical and have low ellipticities, roughly one half of proto-haloes should not collapse in the standard approach to ellipsoidal collapse. This is a strong indication that carefully accounting for proto-halo shapes (and ultimately their boundaries) is key to realistic excursion set modeling.  Another possibility is to attach potential proto-halo shapes to points in the smoothed linear density field by exploiting the correlations between the ratios $a_i/a_j$ and $\\lambda_i/\\lambda_j$ of the eigenvalues of the inertia and tidal tensors (see Figure~\\ref{Fig:inert_deform}). After some experimentation we found a relatively well-defined and mass-independent trend between the ratio $a_3/a_1$ and $\\lambda_3/\\lambda_1$; we show the median trends and their associated one-sigma scatter for several different mass bins in Figure~\\ref{Fig:a3a1_over_e3e1}. Note that the haloes selected for this plot span more than four orders of magnitude in mass, yet at any given mass scale the average minor-to-major axis ratios, $a_3/a_1$, are highly correlated with the relative compression factors along the minor and major axes of the tidal field. The main trend can be reasonably well described by the simple function \\begin{equation} \\frac{a_3}{a_1}=C\\times\\tanh\\biggr(\\frac{\\lambda_3}{\\lambda_1}\\biggl) + \\psi, \\label{eq:a3a1_over_e3e1} \\end{equation} where $C=0.465$ and $\\psi=0.488$ are the best-fit parameters. Eq.~\\ref{eq:a3a1_over_e3e1} is shown as a dotted-dashed line in Fig.~\\ref{Fig:a3a1_over_e3e1}  Given eq.~\\ref{eq:a3a1_over_e3e1}, one can estimate the axis ratios $a_3/a_1$ of potential proto-haloes completely in terms of the tidal field measured at a given point, without making any explicit reference to mass scale. Given an estimate of $a_3/a_1$ we can set $a_2/a_1=\\sqrt{a_3/a_1}$ (i.e. assume perfectly triaxial proto-haloes) and compute the shape-dependent collapse boundaries on a point by point basis with knowledge of only the eigenvalues of the linear tidal field: $\\lambda_1\\geq\\lambda_2\\geq\\lambda_3$. The collapse barrier can then be compared to the true overdensity at that point in the same sense as before.  In Figure~\\ref{Fig:FracAbovee3e1} we plot the fraction of proto-haloes whose average overdensities exceed the collapse threshold inferred in this way. Haloes are binned according to the ratio $\\lambda_3/\\lambda_1$ measured at their centers of mass and the appropriate collapse threshold is then computed from eq.~\\ref{eq:a3a1_over_e3e1}. The curves shown in Figure~\\ref{Fig:FracAbovee3e1} adopt the same colors and line-styles as those in Figure~\\ref{Fig:FracAboveMass}. Solid curves with circles correspond to the fraction of proto-haloes with measured density contrasts exceeding the barrier height inferred from our model; dashed curves to the fraction of haloes with $\\delta\\geq\\delta_{ec}$, the \\citet{Sheth2001} model (eq.~\\ref{eq:ecbarrier}). Clearly, more careful modeling of the intrinsic shapes of dark matter proto-halos can better account for the large scatter in their measured density contrasts.  \\begin{figure} \\begin{center} \\resizebox{8.5cm}{!}{\\includegraphics{deform_vs_inert_MassDep.ps}} \\end{center} \\caption{Median values of axis ratio $a_3/a_1$ for proto-haloes computed in bins of $\\lambda_3/\\lambda_1$. Different colored curves and symbols denote different mass bins, as indicated in the legend. The total number of haloes in each bin is also given. The error bars correspond to the 25$^{th}$ and 75$^{th}$ percentiles. There is a strong and relatively mass independent trend between $a_3/a_1$, which can be approximated by eq.~\\ref{eq:a3a1_over_e3e1}, as shown with the triple-dot-dashed line.} \\label{Fig:a3a1_over_e3e1} \\end{figure}  \\begin{figure} \\begin{center} \\resizebox{8.5cm}{!}{\\includegraphics{frac_above_e3e1_thresh.ps}} \\end{center} \\caption{The fraction of haloes with smoothed, linearly extrapolated overdensities above the threshold predicted by the ellipsoidal collapse model. Fractions are computed in bins of $\\lambda_3/\\lambda_1$, which measures the relative compression factors along the minor and major axes of the linear tidal field. Boundaries are computed from the ellipsoidal collapse model described in Section~\\ref{sec:collapse}, with proto-halo shapes estimated from eq. ~\\ref{eq:a3a1_over_e3e1}. In the model, all proto-haloes are assumed to be perfectly triaxial, and we set $p=0$ for simplicity. The solid curves plot the fraction of haloes with measured density contrasts exceeding the collapse threshold inferred from our model; dashed curves plot the fraction of haloes exceeding the model of \\citet{Sheth2001}. Different colored curves show results for our two simulations; red lines for our $1200\\Mpch$ box simulation; blue for our $150\\Mpch$ box run.} \\label{Fig:FracAbovee3e1} \\end{figure}    We have used two cosmological simulations of structure formation to test a key assumption commonly used in excursion set modeling of structure formation, namely, that the regions of the linear density field that collapse to form bound objects are spherical. We used our results to motivate a deeper look into the collapse boundaries of dark matter halos predicted by the ellipsoidal collapse model. Our main results can be summarized as follows.  \\begin{itemize}  \\item The shapes of the Lagrangian patches that eventually collapse to form CDM halos are highly non-spherical. Over the mass scale $\\sim 10^{11} h^{-1} {\\rm M}_{\\odot}$ to $\\sim 10^{15} h^{-1} {\\rm M}_{\\odot}$ the average intermediate-to-major axis ratio varies from $\\langle a_2/a_1 \\rangle\\sim 0.7$ to $\\sim 0.8$, and $\\langle a_3/a_1 \\rangle$ from roughly $\\sim 0.4$ to $\\sim 0.65$. Above the characteristic mass scale for collapse (roughly $10^{13} h^{-1} {\\rm M}_{\\odot}$) proto-haloes are, on average, ``perfectly triaxial'', with axis ratios $a_2/a_1\\approx a_3/a_2$.  \\item The shapes of dark matter proto-haloes are primarily determined by the relative compression factors associated with the linear tidal field acting upon them (see Fig~\\ref{Fig:inert_deform}). For example, haloes that form in regions of the linear density field where $\\lambda_1$ is large and positive, but $\\lambda_3$ is large and negative, tend to accrete material from very flattened and elongated regions. This is because material is compressed along one axis of the tidal field, and dilated along an orthogonal direction. The net effect of this push-and-pull of material is a strong alignment between the principle axes of the inertia tensors of proto-haloes and those of the tidal field measured at their geometric centers \\citep{Porciani2002b,Lee2000}. The origin of this correlation is rooted in the tidal field, which sets the boundaries of regions of the linear density field that are able to collapse by a given time.  \\item The alignment between the eigenvectors of the linear tidal field and the inertia tensor of proto-haloes can be exploited to simplify the application of ellipsoidal collapse to triaxial density perturbations. In Section~\\ref{sec:collapse} we modified the standard ellipsoidal collapse model of \\citet{Bond1991} in order to account for the non-spherical nature of the Lagrangian patches that collapse to form halos. Relaxing the assumption that these regions are initially spherical yields very different predictions for the thresholds for non-linear structure formation.  \\item In the absence of external tidal forces any non-spherical density perturbation requires a {\\em higher} density contrast for collapse than the corresponding spherical perturbation. This is because the internal tidal forces suppress collapse along the major axis of the perturbation, and correspondingly higher densities are needed to counter this. When external tidal forces are introduced as described in Section~\\ref{sec:collapse}, there is a well-defined minimum barrier height for any non-spherical perturbation; for a given tidal prolateness, $p$, this minimum occurs when the density contrast for collapse along the major and minor axis of the density perturbation are equal (See Fig~\\ref{Fig:BoundaryCompare}).  \\item We provided two possible methods for estimating the typical shape of a proto-halo that may form at any given location in the smoothed linear density field. The first involves measuring the average shapes of proto-haloes as a function of their mass, and using these shapes to compute explicitly the average barrier height for gravitational collapse as a function of halo mass. Another possibility is to estimate directly the proto-halo shape by exploiting the correlation between the axis ratio $a_3/a_1$ and the relative tidal compression factors $\\lambda_3/\\lambda_1$ (see Fig.~\\ref{Fig:a3a1_over_e3e1}). Since proto-halos are roughly triaxial ($a_3/a_2=a_2/a_1$) one can compute the typical shape of a proto-halo that will form at a given location based solely on the ratio $\\lambda_3/\\lambda_1$ measured at that location. As many as $\\simgt 90\\%$ of ($N_{\\rm FOF}\\geq 500$) dark matter proto-haloes have measured linear density contrasts that exceed the collapse thresholds inferred in these ways.  \\end{itemize}  All in all, our results indicate that the standard ellipsoidal collapse model of \\citet{Bond1991} provides a rather incomplete census of the possible collapse thresholds for the formation of CDM halos from Gaussian random fields. In addition to the specifics of the linear tidal field at a given location, the collapse threshold depends on the assumed shape of the Lagrangian region that eventually collapses to form a bound object. This is because the linear overdensity required for collapse in a given tidal field is sensitive to the initial displacement of a given mass element from the perturbation center, thus changing how the collapse barriers depend on auxiliary properties of the tidal field. These results will have important implications for excursion set modeling of the halo mass function, spatial clustering and mass accretion histories, and may provide new insight into the complex process of dark matter halo formation."
    },
    "1107/1107.2934_arXiv.txt": {
        "abstract": "The Baryonic Tully-Fisher Relation (BTFR) is an empirical relation between baryonic mass and rotation velocity in disk galaxies. It provides tests of galaxy formation models in \\LCDM\\ and of alternative theories like MOND.  Observations of gas rich galaxies provide a measure of the slope and normalization of the BTFR that is more accurate (if less precise) than that provided by star dominated spirals, as their masses are insensitive to the details of stellar population modeling. Recent independent data for such galaxies are consistent with $M_b = AV_f^4$ with $A = 47\\pm6\\;\\Aunits$. This is equivalent to MOND with $a_0 = 1.3 \\pm 0.3\\;\\textrm{\\AA}\\,\\textrm{s}^{-2}$. The scatter in the data is consistent with being due entirely to observational uncertainties. It is unclear why the physics of galaxy formation in \\LCDM\\ happens to pick out the relation predicted by MOND. We introduce a feedback efficacy parameter $\\mathcal{E}$ to relate halo properties to those of the galaxies they host.  $\\mathcal{E}$ correlates with star formation rate and gas fraction in the sense that galaxies that have experienced the least star formation have been most impacted by feedback. ",
        "introduction": "The Tully-Fisher relation \\citep{TForig} is one of the strongest empirical correlations in extragalactic astronomy. In addition to being a useful distance indicator \\citep[e.g.,][]{sakai}, the physical basis of this relation is of fundamental importance.  In the context of \\LCDM, the Tully-Fisher relation places tight constraints on galaxy formation models \\citep[e.g.,][]{EL96,MdB98a,MMW98,CR,SN,CR,vdB00, BullockTF,MoMao,MayerMoore,gnedinTF,fabioTF,TGKPR,DRTP,TMZBTDS}. In a broader context, it provides a useful test of theories that seek to modify gravity in order to obviate the need for dark matter \\citep[e.g.,][]{milgrom83,FLAG,weyl,MoG}. These should make specific and testable predictions for the Tully-Fisher relation.  The Tully-Fisher relation was originally posed as an empirical relation between optical luminosity and the width of the 21 cm line.  These are presumably proxies for more physical quantities. The 21 cm line-width is a proxy for rotation velocity \\citep[e.g.,][]{verhTF,YSTF} which is widely presumed to be set by the dark matter halo.  Luminosity is a proxy for stellar mass, which is presumably related to the total mass of the system.  Thus a relation like Tully-Fisher seems natural, with the mass of the parent halo specifying both luminosity and rotation velocity.  Luminosity is not a perfect proxy for mass.  The stellar mass-to-light ratio can, and presumably does, vary with galaxy type.  This leads to Tully-Fisher relations with different slopes and scatter depending on the band-pass used to measure the luminosity \\citep[e.g.,][]{TFdust,verhTF,Cscale,noordTF,mastersTF}. Moreover, low luminosity galaxies tend to fall below the extrapolation of the fit to bright galaxies \\citep{PSnonlinear,MvDG98}.  Yet all these observed relations presumably stem from a single underlying physical relation.  \\citet{FreemanLSB} and \\citet{btforig} suggested that baryonic mass is a more fundamental quantity than luminosity.  For any given stellar population, different mass-to-light ratios will be produced in different bands by the same mass of stars.  More importantly, physics should not distinguish between mass in stars and mass in other forms.  Adding the observed gas mass to the stellar mass results in a Baryonic Tully-Fisher Relation (BTFR) that is linear (in log space) over many decades in mass \\citep{verhTF,BdJ,gurov04,M05,pfennBTF,begum,stark,trach,gurov}. The BTFR appears to be the fundamental physical relation underpinning the empirical Tully-Fisher relation. As such, the true form of the BTFR is of obvious importance.  In this paper, I calibrate the BTFR with gas rich galaxies. The masses of gas rich galaxies can be more accurately measured than those of the star dominated galaxies that traditionally define the Tully-Fisher relation because they are less affected by the systematic uncertainty in the stellar mass-to-light ratio \\citep{btforig,stark,newPRL}. This results in a purely empirical calibration with a minimum of assumptions.  In \\S 2, I discuss the data required construct the BTFR, assemble samples of gas rich galaxies, explore the statistics of the data, and calibrate the relation.  In \\S 3 I examine the implications of the relation for \\LCDM\\ and MOND. I define a feedback efficacy parameter $\\mathcal{E}$ to quantify the mapping between observations and theoretical dark halo parameters.  I estimate the value of the MOND acceleration constant $a_0$ from the empirical BTFR. A summary of the conclusions are given in \\S 4. ",
        "conclusions": "The baryonic mass--velocity relation provides an important test of \\LCDM\\ and alternatives like MOND. I have assembled recent data for gas rich rotating galaxies in order to test these theories. With $M_g > M_*$, the location of these objects in the baryonic mass--rotation velocity plane is effectively measured directly from observations without the usual systematic uncertainty in the stellar mass-to-light ratio.  A long standing prediction \\citep{milgrom83} of MOND is that rotating galaxies will fall on a single mass--velocity relation with slope 4:  $M_b \\propto V_f^4$.  This prediction is realized in multiple independent data sets.  Gas rich galaxies fall where predicted by MOND with no free parameters. There are not many predictions in extragalactic astronomy that fare so well over a quarter century after their publication.  In \\LCDM, the naive expectation of a of Tully-Fisher type relation with a slope around 3 fails to predict the location of gas rich galaxies in the plane of the BTFR. Only the most recent detailed numerical models \\citep[e.g.,][]{fabio,piontek} succeed in reproducing the observed phenomenology.  This is certainly progress.  However, a physical understanding for why galaxy formation in the context of \\LCDM\\ should pick out the particular phenomenology predicted \\textit{a priori} by MOND remains wanting."
    },
    "1107/1107.5039_arXiv.txt": {
        "abstract": "We present a 50\\,ks \\chandra\\ ACIS-I X-ray observation of the \\citeauthor{bower} VLA archival field. The observations reach a limiting sensitivity of $\\sim 10^{-4}$~counts~s$^{-1}$, corresponding to a flux of a few times $10^{-15}$\\,\\esc\\ for the models we explore. The \\chandra\\ observations were undertaken to search for X-ray counterparts to the eight transient sources without optical counterparts, and the two transient sources with optical counterparts seen by \\citeauthor{bower} Neither of the sources with optical counterparts was detected in X-rays. One of the eight optical non-detections is associated with a marginal ($2.4\\sigma$) X-ray detection in our \\chandra\\ image. A second optically-undetected \\citeauthor{bower} transient may be associated with a $z = 1.29$ X-ray detected quasar or its host galaxy, or alternatively is undetected in X-rays and is a chance association with the nearby X-ray source. The X-ray flux upper limits, and the one marginal detection, are consistent with the interpretation of \\citeauthor{ofek:10} that the optically-undetected radio transients are flares from isolated old Galactic neutron stars. The marginal X-ray detection has a hardness ratio which implies a temperature too high for a simple one-temperature neutron star model, but plausible multi-component fits are not excluded, and in any case the marginal X-ray detection may be due to cosmic rays or particle background. The X-ray flux upper limits are also consistent with flare star progenitors at $\\gtrsim 1$\\,kpc (which would require the radio luminosity of the transient to be unusually high for such an object) or less extreme flares from brown dwarfs at distances of around 100\\,pc. ",
        "introduction": "In recent years, authors such as \\citet{bower}, hereafter B07; \\citet{becker:10}; \\citet{bower:11}; \\citet{bannister:11}; \\cite{bell:11}; and others, have made use of archival data (observations of calibration fields, or in some cases, entire telescope archives) to search for radio transients. Other studies have targeted particular fields with observations dedicated to the search for transients \\citep[][and references therein]{ofek:11}.  Such studies are preparing the groundwork for a new generation of high-throughput radio telescopes: both dedicated new observatories --- for example, the Allen Telescope Array \\citep{welch}, LOFAR \\citep{lofar}, and ASKAP \\citep{askap} --- and upgrades to older instruments --- for example, Apertif \\citep{apertif} and EVLA \\citep{evla}. A major goal for these new facilities is to survey large areas of sky at high sensitivity, both to build up deep, wide-field images, and to search for transient and variable sources from a range of progenitors \\citep{lazio}.  Currently, large areas of rate versus sensitivity parameter space remain to be explored \\citep{pigss,paperii}, and the radio transient population is not well understood. In observations where radio transients are detected, sometimes counterparts are seen in existing surveys at other wavelengths. For example, transients may coincide with the nuclear regions of galaxies, in which case (depending on the galaxy redshift and hence the inferred radio luminosity of the transients) they are most likely due to active galactic nuclei (AGNs), or perhaps supernovae (SNe) or gamma ray bursts (GRBs). Or they may be offset from the peak of the optical emission, as is the case for two of the transients seen by B07, leading to their interpretation as SNe or GRBs.  For some radio transients, however, no counterparts are seen in archival images at other wavelengths, and no follow-up observations taken close in time to when the transients were seen are available. Follow-up observations will become more common as radio observatories move increasingly into a regime of issuing transient alerts for rapid follow-up by other facilities, but the field is still in its infancy compared to optical transient facilities which are now issuing transient alerts as a matter of routine \\citep[\\eg,][]{ptf}. However, the lack of galaxy counterparts in deep optical images can help exclude some models for these sources (such as AGNs, SNe and GRBs, unless these sources are at very high redshift or extremely obscured), and the lack of stellar counterparts can also constrain models with Galactic progenitors for these events. B07 consider radio SNe, GRBs, late-type stars, soft gamma-ray repeaters, X-ray binaries, pulsars, microlensing, and reflected solar flares as possible progenitors. They conclude that late-type stars are the most likely progenitors for transients with mJy radio fluxes and no optical counterparts.  \\citet{ofek:10}, hereafter O10, suggest old isolated Galactic neutron stars as a possible progenitor for the sources without optical identifications, and \\citet{nakar:11} suggest that the sources with optical counterparts may be due to neutron star -- neutron star or neutron star -- black hole mergers. The latter are of particular interest because they are expected to also be producers of gravitational waves. X-ray observations can be a powerful discriminant among progenitor models, but information on the X-ray properties of the transients and transient hosts is limited.  With this in mind, we obtained \\chandra\\ observations of the B07 field.  We assume an \\lcdmparm\\ cosmology \\citep{wmap}. Magnitudes are given in the Vega system unless stated otherwise. ",
        "conclusions": "Of the ten B07 transients, eight are undetected in X-rays; one is associated with an X-ray source of marginal ($2.4\\sigma$) significance, which may be due to cosmic rays or particle background; and one may be associated with a QSO or its host galaxy at $z = 1.29$, or may be a spurious positional coincidence. Because of the uncertainty regarding these two associations, it is plausible that all of the B07 transients remain undetected in our X-ray images.  The measured X-ray flux for the transient with a marginal X-ray detection (X124/8L1), and the upper limits for the undetected transients, are consistent with the X-ray luminosity expected for brown dwarfs at distances $\\sim 100$\\,pc, one of the possible progenitors discussed by O10. The small offset between the X124 and 8L1 positions could plausibly be due to the proper motion of such an object.  If instead the transients are due to flaring M stars, most of the radio transients ought to have been detected in quiescence in both the X-ray and OIR images, unless the flares were extreme events (not out of the question, since we know the radio flux densities increased by a factor $\\gtrsim 100$) from progenitors at distances $\\gtrsim 1$\\,kpc.  The X-ray fluxes are also consistent with the suggestion of O10 that the transients could be isolated neutron stars. The radio and OIR fluxes are consistent with this picture as well. If X124 is truly a counterpart to 8L1, hardness ratio constraints appear to rule out a simple one-component spectral model, but there are insufficient data to attempt a more sophisticated fit with additional thermal or non-thermal components.  It is possible that the B07 transients are a mixture of some of the above classes of object, or some of them may be due to some unknown class of transient.  If the progenitors have X-ray fluxes slightly fainter than our sensitivity limit, deeper observations would be expected to detect more of the transients in X-rays, as well as improving the significance of the fainter detection (or rejecting it as a real source). X-ray or radio monitoring campaigns lasting for several months would have a good chance of catching these sources in outburst since for the most likely classes of progenitor, flares are likely to repeat. Near- or mid-infrared data would also be useful to confirm or refute the brown dwarf hypothesis."
    },
    "1107/1107.2275_arXiv.txt": {
        "abstract": "Stellar mass loss rates are an important input ingredient for stellar evolution models since they determine stellar evolution parameters such as stellar spin-down and increase in stellar luminosity through the lifetime of a star. Due to the lack of direct observations of stellar winds from Sun-like stars stellar X-ray luminosity and stellar level of activity are commonly used as a proxy for estimating stellar mass loss rates. However, such an intuitive activity --- mass loss rate relation is not well defined. In this paper, I study the mass loss rate of the Sun as a function of its activity level. I compare in-situ solar wind measurements with the solar activity level represented by the solar X-ray flux. I find no clear dependency of the solar mass flux on solar X-ray flux. Instead, the solar mass loss rate is scattered around an average value of $2\\cdot10^{-14}\\;M_\\odot\\;yr^{-1}$. This independency of the mass loss rate on level of activity can be explained by the fact that the activity level is governed by the large modulations in the solar close magnetic flux, while the mass loss rate is governed by the rather constant open magnetic flux. I derive a simple expression for stellar mass loss rates as a function of the stellar ambient weak magnetic field, the stellar radius, the stellar escape velocity, and the average height of the Alfv\\'en surface. This expression predicts stellar mass loss rates of $10^{-15}-10^{-12}\\;M_\\odot\\;yr^{-1}$ for Sun-like stars. ",
        "introduction": "\\label{sec:Intro}  \\cite{Parker58} introduced the concept of a radially uniform supersonic flow, continuously flowing from an iso-thermal hot corona as a result of a hydrodynamic thermal acceleration. One can imagine space as a huge vacuum cleaner that sucks up solar material, while solar gravity opposes this force. The gas slowly accelerates sub-sonically until it reaches a critical point, where its thermal energy equals the gravitational energy of the Sun. Above this critical point, the flow becomes super-sonic and continues to accelerate to some asymptotic terminal speed, $u_\\infty$. This idea was controversial at first, but have been widely accepted since the first sporadic \\citep{Gringauz60} and consistent \\citep{Neugebauer66} in-situ measurements of the solar wind were made, proving the existence of such flow. Parker's model has served as the baseline to all models for the solar wind and stellar winds; the main difference between the models is the source of the energy accelerating the wind. These models include non iso-thermal winds, polytropic winds, sound waves driven winds, dust driven winds, and line-drive winds \\citep[see][with references therein]{lamerscassinelli99}.  Parker's model is hydrodynamic: the flow is assumed to be parallel to the magnetic field, and interaction between the flow and the fields is neglected. However, stellar magnetic fields play a significant role in the physics of stellar winds since they govern the transport of the ionized gas in the corona \\citep{lamerscassinelli99}. \\cite{Parker58} did not include the solar magnetic field in his wind model, but he discussed extensively the consequences of the coupling between the wind and the field to conclude the structure of the magnetic field in the interplanetary space. \\cite{weberdavis67} have shown that a star can lose angular momentum due to the fact that open magnetic field lines in stellar coronae (along where the wind is flowing) are attached to the star at one end, but are free at the other end beyond the Alfv\\'en point, $r_A$, at which the flow speed equals the Alfv\\'en speed, $u_A=B/\\sqrt{4\\pi\\rho}$. Here $B$ is the magnetic field strength and $\\rho$ is the plasma mass density. Coronal plasma is co-rotating with the star, but it is also frozen in to the magnetic field. Therefore, a torque is applied on the star by an arm with a length of $r_A$. This ``magnetic braking'' explains how stars spin down, where the angular momentum loss rate depends on the stellar mass loss rate. Therefore, stellar mass loss rates and angular momentum loss rates are crucial inputs for stellar evolution models, which then provide input to other models.  Unfortunately, neither stellar winds nor stellar mass loss rates can be directly observed for low-mass stars \\citep[see][and references therein]{gudel07}. Some attempts to estimate stellar mass loss rates were done using chemical separation and H$\\alpha$ profiles \\citep{Michaud86,LanzCatala92,Bertin95}, radio observations \\citep{Abbott80,Cohen82,Hollis85,Lim96,Gaidos00}, and observations of x-ray emission due to charge exchange \\citep{Wargelin01}. These calculations have estimated stellar mass loss rates between $10^{-14}-10^{-10}\\;M_\\odot\\;yr$.  Perhaps the most extensive estimation of stellar mass loss rates was done using measurements of Ly$\\alpha$ absorptions at the edge of stellar heliospheres, where the stellar wind collides with the Interstellar Medium (ISM) \\citep{Wood02,Wood05,Wood06}. The formation of a shock at this location causes the build up of a Hydrogen ``wall'' that can be measured in the Ly$\\alpha$ line. By estimating the ISM properties, the stellar mass flux carried by the wind can be determined. Based on many observations of different stellar types, \\cite{Wood05} derived a scaling law for stellar mass loss rate as a function of x-ray luminosity, which can also be considered as a function of stellar age and activity level. They found that young, active stars (such as the young Sun) have mass loss rate of about $10^{-12}\\;M_\\odot\\;Yr^{-1}$, which is $\\sim$100 times higher than the approximated solar mass loss rate of $2\\cdot10^{-14}\\;M_\\odot\\;Yr^{-1}$, derived from typical solar wind parameters near Earth.  The results in \\cite{Wood05} are presented cautiously, and the assumptions made are clearly stated. However, their intuitive scaling law, which states that the more active the star, the higher the mass flux, is commonly used in an {\\it ad-hoc} manner, despite of the fact that it might be overestimated. Other calculations for the young Sun revealed a mass loss rate which is only 10 times higher than the current value \\citep{HolzwarthJardine07,Cohen09,Sterenborg11}. In fact, the scaling law in \\citep{Wood05} breaks down for the most active stars. In contrast, \\cite{Fisher98} and \\cite{Pevtsov03} used solar data to derive a scaling law between X-ray spectral radiance and magnetic flux. They found a nearly linear relation that holds for many orders of magnitude. Therefore, it is unclear how stellar winds and thus mass flux correlate with stellar magnetic activity and observed X-ray flux.  Unlike the sparse information about stellar winds, the solar wind has been widely observed both remotely and in-situ since the beginning of the space era. These observations, together with extended monitoring of the solar magnetic field, have revealed a much more complex physics, with the solar magnetic field governing the power, the acceleration, and the topology of the solar wind. Magnetic energization of the wind is necessary to explain the discrepancy between the fastest wind predicted by Parker's model (about $600\\;km\\;s^{-1}$) and the observed fast solar wind (about $800-900\\;km\\;s^{-1}$). The exact mechanism in which the magnetic energy is converted to kinetic energy is not fully understood yet. Models to explain this energy conversion consider the effect of wave dumping due to electron/ion resonance frequencies, dumping of Alfv\\'en waves, magnetic energy turbulent cascade, and energy dissipation due to magnetic reconnection \\citep[see][for detailed description and references therein]{Aschwanden05,Cranmer10}. These models are tightly related to the solar corona heating problem. Great progress is currently being made in solving the coronal heating and solar wind acceleration problems due to new observations from the HINODE\\footnote {{\\tt http://www.nasa.gov/mission\\_pages/hinode/index.html}} and SDO\\footnote {{\\tt http://sdo.gsfc.nasa.gov}} missions.  Based on the long-term in-situ measurements of the solar wind taken at 0.3~AU (HELIOS), the outer solar system (Voyager), high heliographic latitudes (Ulysses), as well as near Earth (WIND, ACE), some non-trivial properties of the solar wind have been revealed. These properties are well summarized in \\cite{McComas07} (with references therein). A most notable feature is the bi-modal structure of the solar wind, which seems to be composed of two different populations. The first population is the so-called ``fast wind''. It originates from coronal holes -- the large-scale regions of open field lines (polar regions during solar minimum) -- and it is rather steady, with speeds of the order of $750\\;km\\;s^{-1}$ or more. The second population is the so-called ``slow wind''. This wind is much more sporadic than the fast wind and it is assumed to originate from the vicinity of the large, low latitude helmet streamers during solar minimum. During solar maximum, the slow wind is observed to originate from higher latitudes, as well as from the vicinity of active regions. The two populations are not distinguished by flow speed only, but also by densities, temperatures, composition, abundances, distribution functions, and First Ionization Potential (FIP) levels. While the acceleration of the fast solar wind could, in principle, follow Parker's model, with the contribution of magnetic energy by one of the proposed mechanisms above, the slow solar wind introduces a greater challenge in explaining its acceleration mechanism. An observed inverse correlation between the flow speed and electron temperature, and a strong ion temperature anisotropy with very hot minor species at the low corona are inconsistent with the theory of thermal acceleration. In general, it seems that both populations are at first confined in closed loops, then are later released along open field lines. The fast wind is initially confined near coronal holes, where loops are small, cool, less dense, and short-lived, while the slow wind is confined within the large, hot, dense loops, with the plasma spending more time in the loop before it is released \\citep{Feldman99}. \\cite{Fisk99} showed how the mass flux and Poynting flux at the flux-tube base together with the lifetime of the loops are sufficient to explain the fast solar wind and determining its final speed.  The difference in final speed suggests that the two populations gain different amounts of energy through the acceleration process. This can be generally explained by the different energy per unit mass the populations have (due to their different densities). In addition, due to its fast release the fast wind has more time to gain electromagnetic energy compared to the slow wind. The frequent ejection of the fast wind also explains why it is quasi-steady compared to the sporadic slow wind. Alternatively, the two populations can gain the same amount of energy at the coronal base but have different amount of energy losses in the acceleration process. \\cite{SchwadronMccomas03} proposed that both populations gain similar electromagnetic energy to have high speed values ($\\sim 800\\;km\\;s^{-1}$). The difference in the final kinetic energy results from the wind radiating energy in the form of heat conduction and the advection of thermal energy as it is accelerated, where these radiative terms mostly depend on the maximum temperature of the loops. The loops can be thought as a reservoir of energy that is being transferred to the wind once it is released along open field lines. Since the slow wind is released from large, long-lived, hot loops, it radiates much more then the fast wind, so that its final kinetic energy is smaller. They also showed that their scaling law and its dependency on the maximum loop temperature explains the observed anti-correlation between the electron temperature and final wind speed, the FIP effect \\citep{Schwadron99}, and the anti-correlation between the magnetic flux-tube expansion factor and the wind speed \\citep{wangsheeley90}.  \\cite{SchwadronMccomas06}, and \\cite{SchwadronMccomas08} further showed that the energy function presented in \\cite{SchwadronMccomas03} can be modified to provide a linear correlation between the solar X-ray power and the {\\it open} magnetic flux at the base of the flux-tube. They showed that solar wind data obey this linear relation, that it can be extended to many orders of magnitude, and that it is similar to the relation presented in \\cite{Pevtsov03}. Recently, \\cite{Wang10} have estimated the mass and energy flux at the coronal base using solar wind data taken by ACE and Ulysses in combination with solar magnetic field (magnetogram) data. Based on conservation laws along a magnetic flux tube, he derived an identical formula as in \\cite{SchwadronMccomas06}. The only difference is that while \\cite{SchwadronMccomas06} assumed a constant value for the wind speed far from the Sun, \\cite{Wang10} has used actual data for this parameter. He found that most of the energy flux at the coronal base is used to lift the wind from the gravitational well, i.e. provides the wind with the escape velocity. He also found that while the magnetic field, the mass flux, and energy flux vary a lot at the coronal base, these parameters are rather constant at 1~AU where the variations are balanced by the geometrical expansion of the magnetic flux tube. Similarly, \\cite{SchwadronMccomas06} explained why the solar X-ray luminosity vary by an order of magnitude, while the observed wind power is rather constant. This is due to the fact that the strong solar X-ray emissions come from the closed, hot loops and the wind power is dominated by the open, rather constant solar magnetic flux \\citep{Wang00} . Both \\cite{SchwadronMccomas06} and \\cite{Wang10} have suggested  that the solar wind fluxes are determined by the large-scale, dipolar component of the solar magnetic field, as the heliosphere indeed seems to be composed by two hemispheres with opposite polarity separated by a single current sheet, which is reversed through the solar cycle \\citep{Smith01,Jones03}.  In this paper, I argue that the X-ray flux observed on stars (or stellar activity in general) cannot be a reliable proxy for its wind mass flux and power. In Section~\\ref{SWdata}, I investigate the dependency of the solar wind mass loss rate on the solar X-ray flux using solar wind data taken at different periods of time and at different locations in the solar system. In Section~\\ref{sec:Discussion}, I discuss these results and derive a relation between stellar mass loss rate, and the stellar magnetic and gravitational parameters. I conclude my findings in Section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  I study the dependency of solar mass loss rate on the level of solar activity and solar X-ray flux in order to test whether X-ray flux can be a good proxy for stellar mass loss rates. I exclude short-term transient and spikes from the solar wind and solar X-ray data, in order to better estimate the ambient values of the solar mass loss rate and solar X-ray flux at a particular period of time. I use solar wind data taken at different position in the solar system during different periods of solar activity.  Observations show that the solar mass loss rate is scattered around the value of $2\\cdot10^{-14}\\;M_\\odot\\;yr^{-1}$ with variations of a factor of 2-5, while the solar X-ray flux is modified by an order of magnitude or even higher through the solar cycle. The scatter of the solar mass loss rate is wider for the slow wind and is more narrow for the fast wind, since the slow wind is more sporadic than the rather steady fast wind. The relative constancy of the solar mass loss rate is due to the fact that it is determined by the rather constant solar open magnetic flux. Solar activity is governed by the closed flux, which changes much more dramatically through the solar cycle.  I derive a simple expression for stellar mass loss rates as a function of the stellar ambient weak magnetic field, the stellar radius, the stellar escape velocity, and the average height of the Alfv\\'en surface. This expression predicts stellar mass loss rates of $10^{-15}-10^{-12}\\;M_\\odot\\;yr^{-1}$ for Sun-like stars."
    },
    "1107/1107.4209_arXiv.txt": {
        "abstract": "We study the spectrum of collective excitations in the inhomogeneous phases in the neutron star inner crust within a superfluid hydrodynamics approach. Our aim is to describe the whole range of wavelengths, from the long-wavelength limit which can be described by macroscopic approaches and which is crucial for the low-energy part of the spectrum, to wavelengths of the order of the dimensions of the Wigner-Seitz cells, corresponding to the modes usually described in microscopic calculations. As an application, we will discuss the contribution of these collective modes to the specific heat of the ``lasagna'' phase in comparison with other known contributions. ",
        "introduction": "Neutron stars are fascinating objects, containing matter under extreme conditions of temperature, density and magnetic field. In order to study these celestial bodies, theoretical modeling has to be confronted with observations. A prominent observable is the thermal evolution of isolated neutron stars. Properties of the crust thereby influence the cooling process mainly during the first 50-100 years, when the crust stays hotter than the core which cools down very efficiently via neutrino emission (see e.g. \\cite{Yakovlev07}). Heat transport in the crust is the key ingredient to explain the afterburst relaxation in X-ray transients, too~\\cite{Shternin07,Brown09}. Concerning the models for the thermal relaxation of the crust, the most important microscopic ingredients are thermal conductivity and heat capacity, and to less extent neutrino emissivities.  Here, as a first application, we will concentrate on the heat capacity.  More details about the evaluation of the heat capacity and a discussion of the usually considered contributions can be found in \\cite{Gnedin01}. In what follows, we will concentrate on the particularly interesting case of the neutron star inner crust.  The core of neutron stars is composed most probably of homogeneous neutron rich matter, whereas the crust contains different inhomogeneous structures. The inner crust is thereby characterized by the transition from a lattice of atomic nuclei in the outer crust to homogeneous matter in the core. Ravenhall et al.~\\cite{Ravenhall83} and Hashimoto et al.~\\cite{Hashimoto84} predicted that this transition passes via more and more deformed nuclei. Starting from an almost spherical shape, they could form rods or slabs immersed in a neutron gas at the different densities. These ``spaghetti'' and ``lasagna'' phases are commonly called the nuclear ``pasta''. At higher densities, even closer to the core, other phases such as neutron-gas bubbles inside the dense matter (``swiss cheese'' phase) etc. are expected. The formation of the different structures strongly depends on the relative strength of the nuclear surface energy, Coulomb energy and bulk energy, such that it depends on the nuclear interaction. This prediction has been confirmed within different models for the nuclear interaction, see, e.g.,~\\cite{Oyamatsu93,Pethick95,Watanabe03,Avancini09}. These evaluations have been performed at zero temperature. It is clear that at some critical temperature the pasta structures will disappear due to thermal excitations. However, this melting temperature is of the order of several MeV (see e.g.~\\cite{Watanabe00,Avancini10}).  Another point is that in neutron stars older than several minutes, matter becomes superfluid. A first evidence for superfluidity in neutron stars has been discussed already in 1969~\\cite{Baym69}, shortly after the discovery of the first pulsars, in connection with the observation of ``glitches''.  Since then much effort has been devoted to the question of superfluidity and superconductivity in neutron star matter, for the inner crust as well as for the homogeneous core, see for example~\\cite{Pethick95}.  There is no consensus on the exact value of the energy gaps $\\Delta$ in the inner crust~\\cite{Lombardo00,Gezerlis10}, but the common agreement is that they are of the order of 1 MeV~\\cite{Chamel08}. A pairing gap much larger than the temperature strongly suppresses the contribution of individual neutrons to the specific heat which is thus very much dependent on the pairing strength~\\cite{Gnedin01,Fortin09}. For moderate and strong pairing, the main contributions to the heat capacity considered so far in the crust are thus electrons and lattice vibrations as well as collective excitations of nuclei. However, the superfluid character of neutron star matter induces collective excitations of the neutron gas, not considered before, which can give an important contribution to the heat capacity in certain regions, see ~\\cite{Aguilera08,Pethick10,Cirigliano11}.  The aim of the present paper is to study these collective excitations in the inner crust employing a superfluid hydrodynamics approach. Naturally, there is a vast literature on hydrodynamics for neutron stars in different contexts, including the effects of superfluidity \\cite{Epstein88,Carter98,Andersson02,Carter05,Gusakov05,Carter06,Carter06b}. Most of these models are dedicated to the study of macroscopic neutron star properties, whereas our main aim is to study the excitation spectrum of the crust on much smaller length scales. In spirit this is similar to Refs.~\\cite{Pethick10,Cirigliano11,Sedrakian96}, where hydrodynamic equations are developed to study the superfluid Goldstone boson and (lattice) phonons in the long wavelength limit. However, we are interested in shorter wavelengths at which effects of the inhomogeneous structure will manifest themselves, too. In this sense, our approach is situated in between the long wavelength limit and the completely microscopic calculations (see e.g.~\\cite{Sandulescu04,Khan05,Fortin09}) employing the Wigner-Seitz approximation~\\cite{Negele71}. In the latter case, the wavelengths are limited to the size of the Wigner-Seitz cell, because of the imposed boundary conditions which do not include the coupling between neighboring cells.  The paper is organized as follows. In Sec.~\\ref{sec:formalism}, we describe the superfluid hydrodynamics approach. We summarize the hydrodynamic equations, discuss the boundary conditions and the microscopic input. In Sec.~\\ref{sec:results}, we show our first results. For simplicity, we restrict ourselves to one-dimensional inhomogeneities (lasagna phase) in this exploratory study. We discuss the spectrum of the collective modes and their contributions to the specific heat. A summary and perspectives of our work are exposed in Section~\\ref{sec:summary}.  Throughout the article, $c$, $\\hbar$, and $k_B$ denote the speed of light, the reduced Planck constant, and the Boltzmann constant, respectively. ",
        "conclusions": "\\label{sec:summary} In this paper, we have presented a formalism of superfluid hydrodynamics to treat density-wave propagation in inhomogeneous pasta-like nuclear structures which appear in the inner crust of neutron stars. To account for the periodicity of the structure, we incorporate the Floquet-Bloch boundary conditions. The idea is somewhere in between the approaches of Refs.~\\cite{Pethick10, Cirigliano11}, considering only the long-wavelength limit, averaging over the microscopic details of the structure, and microscopic calculations of the crust within the Wigner-Seitz approximation \\cite{Khan05,Sandulescu04}, valid for wavelengths smaller than the radius of the Wigner-Seitz cell. Concerning the microscopic input for the nuclear equation of state and the geometry of the structure, we followed the work by Avancini et al.~\\cite{Avancini09}.  Within this approach, we have calculated the excitation spectrum of a periodic structure of parallel slabs, the lasagna phase. We have shown that the structure can indeed induce non-negligible effects on the excitation spectrum. In particular, we found that the sound velocity of the usual acoustic mode depends on the direction of the propagation, and, more surprisingly, that there is a second acoustic mode whose dispersion relation is almost independent of $q_z$. In addition, we found different optical branches, similar to the phonon spectrum of ordinary crystals.  We have calculated the specific heat corresponding to this excitation spectrum and found that its contribution is much more important than that of individual neutrons, which is strongly suppressed due to the superfluid gap $\\Delta$. At temperatures relevant for neutron stars, the main contributions to the specific heat come from the electrons and from the acoustic collective modes. The latter cannot be obtained within the Wigner-Seitz approximation. Due to the curious sound mode whose energy is independent of $q_z$, the specific heat due to the collective modes goes like $T^2$ instead of $T^3$. [With the same arguments, one would predict that in a rod structure (spaghetti phase), the specific heat should be linear in $T$.] However, it is not clear whether this feature survives when the Coulomb interaction, which has been neglected here, will be taken into account.  Of course, in order to treat the complex geometry, we were obliged to make a couple of approximations. Contrary to the microscopic approaches based on the Quasiparticle-Random-Phase-Approximation (QRPA) \\cite{Khan05,Sandulescu04}, we rely on the assumption that the modes can be described hydrodynamically, which implies in particular that the local neutron and proton Fermi surfaces stay spherical at all times. This assumption is justified if all spatial variations are slow compared to the superfluid coherence length and the temporal variations are slow compared to the superfluid gap. Both assumptions are not very well fulfilled. However, we have cited examples where hydrodynamics gives reasonable answers even beyond these very restrictive limits, and we believe that the results are at least qualitatively correct. The most serious limitation of the present work is probably that the Coulomb interaction has been neglected. The Coulomb interaction between the protons results in an additional coupling between neighboring cells, which can have important consequences for the excitation spectrum. In the approaches of Refs.~\\cite{Pethick10, Cirigliano11}, it was accounted for by including the elasticity of the Coulomb lattice. In our more microscopic approach, the Coulomb potential would have to be included from the beginning into the proton chemical potential $\\mu_p(\\vek{r},t)$ in the Euler equation (\\ref{eq:euler}). This is a difficult task which will be left for future studies.  It has to be stressed that the contribution of the collective modes studied here is potentially more important than other contributions, notably the contribution from individual neutrons. Therefore it is interesting to pursue their investigation and to include the additional contribution to the specific heat in studies of neutron star thermal evolution."
    },
    "1107/1107.5105_arXiv.txt": {
        "abstract": "We revisit the concept of a blazar sequence that relates the synchrotron peak frequency ($\\nu_{\\mathrm{peak}}$) in blazars with synchrotron peak luminosity ($\\mathrm{L}_{\\mathrm{peak}}$, in $\\nu\\mathrm{L}_\\nu$) using a large sample of radio-loud AGN. We present observational evidence that the blazar sequence is formed from two populations in the synchrotron $\\nu_{\\mathrm{peak}}$ $-$ $\\mathrm{L}_{\\mathrm{peak}}$ plane, each forming an upper edge to an envelope of progressively misaligned blazars, and connecting to an adjacent group of radio galaxies having jets viewed at much larger angles to the line of sight. When binned by jet kinetic power ($\\mathrm{L}_\\mathrm{kin}$; as measured through a scaling relationship with extended radio power), we find that radio core dominance decreases with decreasing synchrotron $\\mathrm{L}_{\\mathrm{peak}}$, revealing that sources in the envelope are generally more misaligned. We find population-based evidence of velocity gradients in jets at low kinetic powers ($\\sim 10^{42}-10^{44.5}$ erg s$^{-1}$), corresponding to FR I radio galaxies and most BL Lacs.  These low jet power `weak jet' sources, thought to exhibit radiatively inefficient accretion, are distinguished from the population of non-decelerating, low synchrotron-peaking (LSP) blazars and FR II radio galaxies (`strong' jets) which are thought to exhibit radiatively {\\it efficient} accretion. The two-population interpretation explains the apparent contradiction of the existence of highly core-dominated, low-power blazars at both low and high synchrotron peak frequencies, and further implies that most intermediate synchrotron peak (ISP) sources are not intermediate in intrinsic jet power between LSP and high synchrotron-peaking (HSP) sources, but are more misaligned versions of HSP sources with similar jet powers. ",
        "introduction": "\\label{intro}  The success of the current model for radio-loud Active Galactic Nuclei (AGN) as accreting super-massive black holes with bipolar, relativistic jets has rested on the ability of this physical paradigm to successfully explain the chameleon nature of the spectral energy distribution (SED) of this family of extreme objects. The jet model is inherently anisotropic and allows us to unify the highly luminous and variable blazars as relativistically beamed counterparts of Fanaroff-Riley (FR) radio galaxies \\citep{fan74}.  The typical blazar spectrum exhibits a double peak (in $\\nu F_{\\nu}$), attributed to synchrotron and inverse Compton (IC) radiation from the relativistic jet, which dominates in blazars due to the alignment of the jet axis near to our line-of-sight \\cite[for a review of leptonic blazar models, see][]{bot07}. The synchrotron component peaks from sub-infrared (IR) energies up to hard X-rays, with the IC component peaking at much higher energies (typically gamma-rays). Spectroscopically, blazars are traditionally divided between optically `featureless' BL\\,Lac objects and (broad) emission line-emitting Flat Spectrum Radio Quasars (FSRQ). The long-standing orientation-based unification scheme for radio-loud AGN proposes BL Lacs and FSRQ as counterparts to FR I and FR II radio galaxies, respectively, based on similar spectra, morphologies, and range in extended radio luminosity \\citep{urr95}.  However, violations to this scheme are well-known, and findings of powerful, FR II-like BL Lacs, and low-power FSRQ \\citep{lan06,lan08,kha10} appear to break the simplest dichotomy (going back to the original FR classification) between the lobe-dominated, high-power FR II galaxies and the low-power, edge-darkened FR Is. \\cite{kha10} also found BL Lacs displaying SEDs and hot-spots typical of quasars while others have observed BL Lacs exhibiting broad lines in low continuum states \\citep{ver95,cor00,rai07}.     Further, the importance of optical spectral type in connecting radio galaxies with their aligned counterparts is difficult to assess due to the strong effects of relativistic beaming on equivalent width measurements on which the blazar classifications are based \\citep{geo98_thesis,ghi11_bllacs_fsrq}.  At the same time, the designation of Low and High Excitation Radio Galaxies (LERG and HERG, corresponding to prominence of emission lines as measured with an excitation index, see \\citealt{hin79}) is increasingly found to be mixed among FR I and FR II morphological types; all of this creates serious problems for a unification scheme which connects spectral type with the morphology.   \\subsection{The Blazar Sequence}  The synchrotron peak frequency ($\\nu_{\\mathrm{peak}}$) may take on a wide range of values (from 10$^{12}$ to 10$^{18}$ Hz at the extremes), and is one of the principle ways to classify individual blazars. While FSRQ sources are typically found to have low $\\nu_{\\mathrm{peak}}$ ($<$ 10$^{14.5}$ Hz), BL Lacs are found to span the entire range, with low, intermediate, and high-peaking BL Lacs (LBL, IBL, and HBL) typically dominating surveys selecting in the radio, optical, or X-ray, respectively. We adopt the generic terms for low, intermediate, and high {\\sl synchrotron-peaking} (LSP, ISP, HSP) blazars independently of the spectroscopic type, as in \\cite{abd10_LSP_ref}.  \\cite{fos98} found an anti-correlation between the frequency of the synchrotron peak and the bolometric and synchrotron peak luminosities and suggested a link between the power of the source and $\\nu_{\\mathrm{peak}}$. This sequence appeared to match a spectroscopic one, from predominantly powerful FSRQ sources at the low-peak, high-luminosity end through LBLs and finally HBLs, at the low-luminosity end.  \\cite{ghi98} suggested that more efficient cooling of particles in the jets of high luminosity blazars is responsible for the lower peak frequencies. However, it is not clear how the continuous blazar sequence fits in with the bi-modality of morphology and spectral type in radio galaxies.  \\subsection{Beaming and Quantitative Unification} \\label{velocitygrad}  Orientation-based unification predicts that misaligned blazars will drop in luminosity and synchrotron peak frequency according to the decrease in Doppler boosting with increasing angle.  In the simplest case (a single Lorentz factor and a convex spectrum), the ratio of decrease in log frequency to log luminosity will be approximately 1:4, roughly orthogonal to the sequence. \\citep[This holds for moving features in the jet; for a standing shock model the relation is 1:3, see {\\sl e.g.,}][]{lin85}. The connection between blazars and the {\\sl nuclear emission} in radio galaxies can naturally inform our jet models.  The synchrotron nature of the nuclear radio galaxy emission has been confirmed by comparison of the broad-band spectral indices and SED characteristics of FR I nuclei with BL Lacs \\citep{chi99,cap00,chi00,cap02,tru03}, where the effects of relativistic beaming (rather than obscuration) accounted for the 10$-$10$^4$ difference in radio and optical luminosities between the BL Lacs and radio galaxies.  \\cite{cap02} and \\cite{tru03} agree upon jet bulk Lorentz factors of $\\Gamma$ $\\sim$ 4$-$6 in order to account for the difference in radio galaxy and blazar luminosities, while \\cite{har03} found $\\Gamma$ $\\sim$ 3 sufficient to match the luminosity functions of the B2 and EMSS and 1 Jy samples. These values agree with VLBI studies of the pc-scale jets of HSP BL Lacs which show that they are barely super-luminal, with average bulk Lorentz factors $\\Gamma$ $\\sim$ 2$-$3 \\citep{pin01_tev_bllacs,edw02_tev_bllacs,gir04,pin08,pin10}; however, several LSP BL Lacs show high velocities on VLBI scales \\citep{jor01,coh07}. Interestingly, \\cite{har03} also found that they could not reproduce the redshift distribution of 1 Jy BL Lac sample unless some of the high-redshift BL Lacs corresponded to FR IIs.  In contrast, high Doppler factors ($\\sim$ 50 or greater) must be invoked to explain the brightness and rapid variability of TeV blazars, as the jet must be highly relativistic in order to avoid absorption by co-spatial IR photons \\citep{don95,aha07,ghi09_tev_variability,gia09}. The discrepancy in apparent jet speeds and radio galaxy nuclear emission can be resolved at least for FR I/BL Lac sources by invoking the presence of structured jets with multiple emitting regions of varying speeds. This includes the spine-sheath axisymmetric model explored by \\cite{swa98}, \\cite{chi00}, and \\cite{ghi05}, and the decelerating model of \\cite{geo03}, all of which imply a high-speed spine or base which dominates the flux when the source is aligned, and a slower component which dominates if the source is misaligned. Other blazar studies also support the decelerating model for low power blazars, as recent VLBI observations of the MOJAVE sample\\footnote{MOJAVE (Monitoring Of Jets in Active galactic nuclei with VLBA Experiments): Over 250 bright blazars being continuously monitored with VLBI.  (www.physics.purdue.edu/mojave)} reveal that the these sources also have lower apparent jet speeds, versus the high-power FSRQ which maintain collimated jets and consistent jet speeds over longer distances \\citep{lis09}.  \\subsection{The Blazar Envelope} The simplest scenario behind the appearance of a blazar sequence is that the physical properties of extragalactic jets (as measured through the synchrotron peak frequency and luminosity) are a function of a single parameter, the jet kinetic power ($\\mathrm{L}_\\mathrm{kin}$). If there is a one-to-one correspondence between $\\nu_\\mathrm{peak}$ and $\\mathrm{L}_\\mathrm{kin}$, the original blazar sequence could be seen as the extreme `edge' of an envelope, consisting of the most aligned (and therefore brightest) objects.  More misaligned sources should appear below the sequence, causing it to widen.  If the jets are monoparametric in terms of jet power, these misaligned sources should fall in a line away from their aligned counterparts on the sequence, according to the drop in Doppler factor, forming `tracks' of similar kinetic power in the envelope of blazars populating the $\\nu_\\mathrm{peak}$ - $\\mathrm{L}_\\mathrm{peak}$ plane.  The original blazar sequence was based on 126 objects, a few dozen of which were detected in gamma-rays. Over a decade later there are thousands more blazars and blazar candidates in the literature and a much greater density of multi-wavelength data in general. Many of the newer samples have been designed to combat the early selection effects of the radio and X-ray surveys, such as the low radio power {\\it CLASS} sample \\citep{cac04_class_testing_the_sequence}, and it is now clear that as our surveys go deeper, many more low-luminosity blazars are emerging. These newer samples challenge both the sequence and its simplest extension that considers the sources in the envelope simply to be de-beamed blazars. Contrary to what is anticipated by the sequence, \\cite{lan08} showed that HSP and LSP BL Lacs objects cover similar ranges in extended radio power. \\cite{cac04_class_testing_the_sequence} found several low luminosity, low $\\nu_\\mathrm{peak}$ sources with relatively high core dominance compared to brighter objects at similar low $\\nu_\\mathrm{peak}$, suggesting that the former are not simply related to the latter through misalignment, as one might naively expect.  In this paper we examine the available evidence for a `blazar envelope' beneath the upper-limit of the most aligned blazars, due primarily to mis-alignment. We test the possibility that jets are monoparametric engines by examining how the source power is related to other properties of the blazar such as location in the $\\nu_{\\mathrm{peak}}-\\mathrm{L}_{\\mathrm{peak}}$ plane, and re-evaluate the blazar sequence in light of the great increases in multi-wavelength, multi-epoch data available for blazar and radio galaxy jets.    \\begin{deluxetable*}{l l c p{7cm}} \\tablecaption{Blazar Samples\\label{table:sources}} \\tablewidth{0pt} \\tablehead{ \\colhead{Sample} & \\colhead{Abbrev.} & \\colhead{No. of}  & \\colhead{References}\\\\ & & Objects\\tablenotemark{$\\dagger$} & \\\\ } \\startdata 1 Jansky blazar sample & 1 Jy & 34 & \\cite{sti91_1jybll} \\\\ 2-degree field (2dF) QSO survey & 2QZ & 56 & \\cite{lon02_2qz_i,lon07_2qz} \\\\ 2 Jansky survey of flat-spectrum sources& 2 Jy & 213 &  \\cite{wal85} \\\\ The Candidate Gamma-Ray Blazar Survey & CGRaBS & 1625 &  \\cite{hea08_cgrabs} \\\\ Cosmic Lens All-sky Survey & CLASS & 236 & \\cite{mar01_class_i,cac02_class_ii,cac04_class_testing_the_sequence} \\\\ Deep X-ray Radio Blazar Survey & DXRBS & 283 & \\cite{per98_dxrbs_i,lan01_dxrbs_ii,pad07_dxrbs_iii}\\\\ \\emph{EGRET} blazar sample & ... & 555 &  \\cite{muk97_egret_blazars_update} \\\\ \\emph{Einstein} Slew survey sample of BL Lac objects & ... &48 &  \\cite{per96_slewbll} \\\\ \\emph{Fermi} blazar sample & 1FGL & 689 &  \\cite{abd10_fermi_cat} \\\\ Hamburg-RASS bright X-ray AGN sample & HRX & 104 & \\cite{bec03_hrx_evolution} \\\\ Parkes quarter-Jansky flat-spectrum sample & ... & 878 & \\cite{jac02_parkes_i, hoo03_parkes_ii, wal05_parkes_iii} \\\\ Radio-Emitting X-ray source survey & REX & 55 &  \\cite{cac99_rex_i} \\\\ Radio-Optical-X-ray catalog  & ROXA & 816 &  \\cite{tur07_roxa} \\\\ RASS - Green Bank BL Lac sample & RGB & 127 & \\cite{lau99_rgb_sample} \\\\ \\emph{Einstein} Medium-Sensitivity Survey of BL Lacs & EMSS & 44 & \\cite{mor91_emss_bllacs,rec00_emssbll_ii} \\\\ RASS - SDSS flat-spectrum sample & ... & 501 & \\cite{plo08_sf}\\\\ Sedentary survey of high-peak BL Lacs & ... & 150 &  \\cite{gio99_sed_i,gio05_sed_ii, pir07_sed_iii} \\\\ \\enddata \\tablenotetext{$\\dagger$}{Number of sources listed is total sources in sample (not necessarily number of blazars).} \\end{deluxetable*}  The paper is organized as follows: in \\S \\ref{methods} we explain the sample selection and describe our statistical and phenomenological models which allowed us to estimate the important envelope parameters for blazars and radio galaxies. In \\S \\ref{results} we present the results of the first round of data collection and analysis and the `blazar envelope'. In \\S \\ref{discussion} we discuss the implications of the blazar envelope and the potential for constraining jet models. In \\S \\ref{conclusions}, we summarize our findings. Throughout the paper we adopt a cosmology based on the concordance model with $\\Omega_M$ = 0.27, $\\Omega_{\\Lambda}$ = 0.73, and $H_0 = 71 \\rm ~ km ~ s^{-1} ~ Mpc^{-1}$. Energy spectral indices are defined such that $f_\\nu \\propto \\nu^{-\\alpha}$, and all luminosities (L$_\\nu$) are references to powers ($\\nu$L$_\\nu$ implied). ",
        "conclusions": "\\label{conclusions} In this work we have analyzed the multi-frequency spectrum of several hundred blazars in order to accurately measure the synchrotron peak frequency and luminosity, and performed spectral analysis of the low-frequency emission in order to estimate the intrinsic jet power by a scaling of the 300 MHz extended radio power. The connection of blazars and radio galaxies in the same plane amounts to an extension of the original blazar sequence to a blazar envelope, in which highly aligned blazars form an upper limit in the $\\nu_{\\mathrm{peak}}-\\mathrm{L}_{\\mathrm{peak}}$ plane, while radio galaxies populate the lower left. We find that the jet power does not absolutely determine the $\\nu_\\mathrm{peak}$ of a maximally-beamed source. Instead, there seems to a break-down at low kinetic jet power, such that these jets are capable of having either low or high synchrotron peaks when aligned with the line of sight.  Our results suggest a shift away from the continuous blazar sequence interpretation, but confirmation waits for ever more data. We summarize the most important points here:    \\begin{itemize} \\item Among our sample of over 600 blazars with well-sampled blazar SEDs, none with high peak frequencies and high luminosities are found. This void may be explained by the original arguments advanced by \\cite{ghi98}, that higher external radiation fields in high-power sources lead to greater cooling and thus lower break/peak energies.  \\item However, it is not clear that there is a continuum of SED types with jet power. The source power alone does not appear to determine the placement of a source along the aligned sequence.  Instead, we suggest that radio-loud AGN can be divided into two populations: low-synchrotron-peaking sources without significant velocity gradients or deceleration (strong jets) and high-synchrotron-peaking sources which do have such structure (weak jets).  The attributes of these samples are summarized in Table \\ref{table:ov}.  \\item It appears that below a certain jet power, sources can either have strong or weak jets, suggesting another parameter besides kinetic power is responsible for the jet structure. As in \\cite{ghi08_new_perspective}, we suggest that this parameter is the accretion rate $\\dot{m}$ in Eddington units with the two families separated at $\\dot{m}=\\dot{m}_{cr}\\sim 3\\times 10^{-3}-10^{-2}$, with sources having $\\dot{m}<\\dot{m}_{cr}$ exhibiting radiatively inefficient accretion \\citep{narayan97}.   \\item A key prediction of the two-population scenario is that ISP sources primarily consist of misaligned HSP objects, as suggested by their slightly lower $\\mathrm{L}_{\\mathrm{peak}}$ values and similar $\\mathrm{L}_\\mathrm{kin}$ to HSPs. Observations of ISP sources with synchrotron peak values and core dominance values between those typical of the LSP and HSP sources (currently missing) would support a continuous sequence interpretation instead.  \\item Observations of highly core-dominated, low-peaking BL Lacs (such as BL Lac itself), including several BL Lacs exhibiting FR II-like morphology, leads us to speculate that these sources belong to the strong-jet parent population, along with FR II radio galaxies. Likewise, some ISP and HSP sources with quasar-like spectra in our sample belong to the weak population. Thus the spectral types associated with strong and weak jets are mixed. \\item The low peak frequencies seen in our FR I radio galaxies supports the interpretation of a horizontal movement in the $\\nu_{\\mathrm{peak}}-\\mathrm{L}_{\\mathrm{peak}}$ plane for weak sources which supports a model with velocity profiles in the jet and leads to our two-population scheme.  Observations of high synchrotron peaks in radio galaxies would be very difficult to explain under this scenario. \\end{itemize}"
    },
    "1107/1107.2808_arXiv.txt": {
        "abstract": "We have surveyed spatial profiles of the Fe K$\\alpha$ lines in the Galactic center diffuse X-rays (GCDX), including the transient region from the GCDX to the Galactic ridge X-ray emission (GRXE), with the Suzaku satellite. We resolved Fe K$\\alpha$ line complex into three lines of Fe \\emissiontype{I}, Fe \\emissiontype{XXV} and Fe \\emissiontype{XXVI} K$\\alpha$, and obtained their spatial intensity profiles with the resolution of $\\sim \\timeform{0D.1}$. We compared the Fe \\emissiontype{XXV} K$\\alpha$ profile with a stellar mass distribution (SMD) model made from near infrared observations. The intensity profile of Fe \\emissiontype{XXV} K$\\alpha$ is nicely fitted with the SMD model in the GRXE region, while that in the GCDX region shows $3.8\\pm0.3$ $(\\timeform{0D.2}<|l|<\\timeform{1D.5})$ or $19\\pm6$ $(|l|<\\timeform{0D.2})$ times excess over the best-fit SMD model in the GRXE region. Thus Fe \\emissiontype{XXV} K$\\alpha$ in the GCDX is hardly explained by the same origin of the GRXE. In the case of point source origin, a new population with the extremely strong Fe \\emissiontype{XXV} K$\\alpha$ line is required. An alternative possibility is that the majority of the GCDX is truly diffuse optically thin thermal plasma. ",
        "introduction": "% Diffuse X-rays were discovered from the inner disk of the Milky Way Galaxy with the HEAO-1 satellite \\citep{Wo82}, and called as the Galactic ridge X-ray emission (GRXE).  EXOSAT observed more details and found that the GRXE were extended by about $\\pm \\timeform{60D}$ in longitude and $\\pm (\\timeform{1D}$--$\\timeform{2D}$) in latitude \\citep{Wa85}. The Tenma satellite discovered strong emission lines at the energy of about 6.7 keV (here the 6.7-keV line) in the GRXE \\citep{Ko86}. The Ginga satellite made survey observations along the Galactic plane with the 6.7-keV line, and found a flux peak at the Galactic center with the width of about $\\timeform{1D}$. We call this peak as the Galactic center diffuse X-rays (GCDX) \\citep{Ko89, Ya90, Ya93}.  The ASCA satellite separated the 6.7-keV line in the GCDX into three emission lines at 6.4, 6.7 and 7.0 keV, which were respectively identified as K$\\alpha$ lines from neutral (Fe \\emissiontype{I}), He-like (Fe \\emissiontype{XXV}) and H-like (Fe \\emissiontype{XXVI}) iron \\citep{Ko96}. For the GRXE, the Suzaku satellite resolved the Fe K$\\alpha$ lines \\citep{Eb08}. The strongest emission line in the GCDX and GRXE is Fe \\emissiontype{XXV} K$\\alpha$ at 6.7 keV, and hence would be due to an optically thin thermal plasma. From the flux ratio of Fe \\emissiontype{XXVI} K$\\alpha$/Fe \\emissiontype{XXV} K$\\alpha$, the plasma temperature of the GCDX was estimated to be around $kT=$5--10 keV (e.g. \\cite{Ko07b, Ya09}).  Since spectra of the GCDX and GRXE have spatial temperature variations and/or may be contaminated by possible non-thermal emissions, we concentrate on the flux of Fe \\emissiontype{XXV} K$\\alpha$ line as a pure component of the thin thermal plasma. Previous studies of the Fe K$\\alpha$ lines with the Ginga and RXTE satellites \\citep{Ya93,Re06}  had  limited spatial resolutions (FWHM $\\sim \\timeform{1D}$--$\\timeform{2D}$)  compared to the size of the GCDX, and the Fe \\emissiontype{XXV} K$\\alpha$ line flux was contaminated by the adjacent Fe \\emissiontype{I} and \\emissiontype{XXVI} K$\\alpha$ lines due to the limited energy resolutions (FWHM $\\sim$ 1.1 keV at 6 keV). \\citet{Ya09} reported that the intensity ratio of  the three Fe K$\\alpha$ lines has spatial dependence and hence to resolve the three lines is essential for the study the nature of Fe \\emissiontype{XXV} K$\\alpha$ line. This paper reports, for the first time, the Fe \\emissiontype{XXV} K$\\alpha$ line distributions in the GCDX and the transient region from the GCDX to the GRXE ($\\timeform{-3D}< l < \\timeform{2D}$, $\\timeform{-2D} < b <\\timeform{1D}$) with the CCD cameras onboard the Suzaku satellite.  The observations and data reduction are described in section 2. The analysis and the intensities of Fe \\emissiontype{I} K$\\alpha$, Fe \\emissiontype{XXV} K$\\alpha$ and Fe \\emissiontype{XXVI} K$\\alpha$ are given in section 3. In section 4, we compare the Fe \\emissiontype{XXV} K$\\alpha$ profiles of the GCDX and GRXE with that of a stellar mass model made from near infrared (NIR) observations. Based on the profile comparison between NIR stars and Fe \\emissiontype{XXV} K$\\alpha$, we discuss origin of the GCDX. In this paper, uncertainties are quoted at the 90\\% confidence range unless otherwise stated. ",
        "conclusions": "% \\subsection{The distributions of the Fe \\emissiontype{XXV} K$\\alpha$ line and the stellar mass} % Using the intensity profile, origin of the Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX is discussed in comparison with that in the GRXE. We fitted the profile of the Fe \\emissiontype{XXV} K$\\alpha$ line with a stellar mass distribution (SMD) model which was originally compiled  by \\citet{Mu06}, based on the NIR observations \\citep{La02, Ke91}. The SMD model consists of four components of the nuclear stellar cluster (NSC), the nuclear stellar disk (NSD), the Galactic bulge (GB) and the Galactic disk (GD). The details of the SMD model are described in Appendix.  We made simultaneous fit for the longitudinal distribution (figure 1a) and for many latitudinal distributions (figure 1b). In this fitting, we included the results of  \\citet{Ko07b} ($|l_*|<\\timeform{0D.4}$), \\citet{Ya09} $(|l_*| > \\timeform{8D})$ and the result of  $(l_*, b_*) =(\\timeform{0D.056}, \\timeform{-1D.954})$. These data were added in the longitudinal plot in figures 4a and c, except for those of \\citet{Ya09} of $|b| > \\timeform{0.15D}$. The additional data points in figures 4a and c  are not exactly on $b_*=\\timeform{0D}$. However, in the simultaneous fitting, the real coordinate values were taken into account.  \\subsubsection{The GRXE region}\\label{ch:GRXEregion} % The intensity distribution of the Fe K$\\alpha$ line along the Galactic plane has been already reported by \\citet{Ya93} and \\citet{Re06}. Since \\citet{Re06} found that the Fe K$\\alpha$ line intensity profile agrees well with the NIR profile in the GRXE region using the RXTE data, we first tried to fit the data of $|l_*|>\\timeform{1D.5}$ regions and $|b_*|>\\timeform{0D.5}$ regions simultaneously with the SMD model. We assumed that the mass emissivities (see equation A7) for each component are the same, i.e. $\\varepsilon_{\\rm NSC}= \\varepsilon_{\\rm NSD} =\\varepsilon_{\\rm GD}=\\varepsilon_{\\rm GB}=\\varepsilon$. Thus the free parameter was only one, $\\varepsilon$ or the normalization. The best-fit parameter was $\\varepsilon= (5.85\\pm0.03) \\times 10^{33}$ photons s$^{-1}$ $M_{\\solar}^{-1}$. The fitting result for the longitudinal distribution is given in figure 4a, while that of the latitude  at $l_*=\\timeform{-0D.114}$ is given in figure 4b as a typical example.  As is shown in figure 4a, we obtained a nice fit in the GRXE data of $|l_*| > \\timeform{1D.5}$, in which the GD and partly GB components mainly contribute in the SMD model. Figure 4b also shows reasonable fit in the region of $|b_*| > \\timeform{0D.6}$, where the GD and GB components contribute equally.  We checked the consistency between the Suzaku results (figure 4) and those of \\citet{Re06} (figure 3 in their paper). The data in the GRXE region of $|l_*|=\\timeform{10D}$--$\\timeform{20D}$ were compared since the limited spatial resolution of RXTE made it impossible to resolve the detailed structure around the Galactic center (GCDX). The result of figure 4a is purely for  Fe \\emissiontype{XXV} K$\\alpha$ line and the Fe \\emissiontype{XXV}~K$\\alpha$ line intensity at $|l_*|=\\timeform{10D}$--$\\timeform{20D}$ is about $1\\times10^{-7}$ photons cm$^{-2}$ s$^{-1}$ arcmin$^{-2}$, while that of \\citet{Re06} is the sum of Fe \\emissiontype{I}, \\emissiontype{XXV} and \\emissiontype{XXVI} K$\\alpha$ lines of about $4.5 \\times 10^{-4}$ photons cm$^{-2}$ s$^{-1}$ deg $^{-2}$  (hence $1.3\\times10^{-7}$ photons cm$^{-2}$ s$^{-1}$ arcmin$^{-2}$). The Fe K$\\alpha$ line complex in the RXTE data was not separated into the three lines because of the limited energy  resolution. From table 1 of \\citet{Eb08} and also from figure 3 of \\citet{Ya09}, we estimated the mean flux ratio of  Fe \\emissiontype{I}, \\emissiontype{XXV} and \\emissiontype{XXVI} K$\\alpha$ lines as $0.2:1.0:0.2$. Thus the Fe \\emissiontype{XXV} K$\\alpha$ flux is about 70\\% of the total flux of Fe K$\\alpha$ line complex. Taking this flux ratio into account, our results are consistent with that of \\citet{Re06}. We, therefore, do not discuss origin of Fe \\emissiontype{XXV} K$\\alpha$ line in the GRXE further than \\citet{Re06}.   \\subsubsection{The combined regions of the GCDX and GRXE} In the previous section \\ref{ch:GRXEregion}, we found that Fe \\emissiontype{XXV} K$\\alpha$ profile in the GCDX shows clear excess above the SMD model which gave a nice fit to that in  the GRXE. As is shown in figure 4a, this excess in the GCDX is confined in the region of the NSC and NSD. Therefore we fitted the whole data of Fe \\emissiontype{XXV} K$\\alpha$ line (the GCDX+GRXE region) simultaneously, fixing $\\varepsilon_{\\rm GD}=\\varepsilon_{\\rm GB}= 5.85 \\times 10^{33}$ photons s$^{-1}$ $M_{\\solar}^{-1}$ and two free parameters of $\\varepsilon_{\\rm NSC}$ and $\\varepsilon_{\\rm NSD}$. The best-fit results are given in figures 4c and d, which show globally a good fit with the enhanced  parameter values of $\\varepsilon_{\\rm NSC}$ and $\\varepsilon_{\\rm NSD}$. The best-fit $\\varepsilon_{\\rm NSC}$ and $\\varepsilon_{\\rm NSD}$ are listed in table 4. As is found from table 4, we needed $3.8\\pm0.3$ and $19\\pm6$ times higher emissivities in the NCD and NSC region (GCDX) than that in the GD+GB region (GRXE).  The fitting in which the mass emissivities of the NSD and NSC ($\\varepsilon_{\\rm NSD}$ and $\\varepsilon_{\\rm NSC}$) are fixed to the same value of the GD and GB, but the stellar masses of the NSC and NSD  ($\\rho_{\\rm c}$ and $\\rho_{\\rm d}$) are free parameters (see equation \\ref{eq:nsc} and \\ref{eq:nsd})  gives that respective stellar masses of  the NSC and the NSD are $3.8\\pm0.3$ and $19\\pm6$ times larger than those obtained with the NIR observations \\citep{La02}, normalized to the Galactic ridge.  In section 4.3, we will discuss the origin of  Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX based on these results.   \\subsection{Systematic error for Fe \\emissiontype{XXV} K$\\alpha$ line profile and the SMD model} Since the statistics is limited, present spectral analysis in the small region (\\timeform{0.1D}$\\times$\\timeform{0.2D}) cannot confine the interstellar absorption ($N_{\\rm H}$) within reasonable errors. We therefore did not correct the absorption effect for the flux of Fe \\emissiontype{XXV}~K$\\alpha$. The largest interstellar absorption is $N_{\\rm H} \\sim 6 \\times10^{22}$ cm$^{-2}$ near the Galactic center \\citep{Pr95}. Using the cross section of the photoelectric absorption of \\citet{Mo83}, we estimated that the flux of Fe \\emissiontype{XXV} K$\\alpha$ is reduced only by 9\\% in the largest case of $N_{\\rm H}=6 \\times10^{22}$ cm$^{-2}$. Since the largest reduction is expected in the GCDX region, the excess discussed in section 4.1.2 becomes even larger by about 10\\%, if we correct the interstellar absorption. A larger error is in the making of the SMD model,  which is estimated to be about 50\\% \\citep{Mu06} (see Appendix). Taking these errors into account, our conclusion of the large excess of Fe \\emissiontype{XXV}~K$\\alpha$ line at the GCDX to the GRXE in the best-fit SMD model is still valid, and hence no essential change of the discussion is required.   \\subsection{Origin of Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX} We found that  the mass emissivity of Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX region (NSC and NSD) shows excess of $3.8\\pm0.3$ (NSC) and $19\\pm6$ (NSD) times larger than that in the GRXE region (GD and GB), or the stellar masses in the NSC and NSD regions are $3.8\\pm0.3$ and $19\\pm6$ times larger than those estimated  with the NIR observations. Integrating over the whole region of the NSC and NSD, we obtain that the total stellar mass in the NSC and NSD are about $1.4 \\times 10^9$~$M_{\\solar}$ for the former case, which is fully consistent with the kinematic mass distribution (see figure 14 of \\cite{La02}). However the latter case  requires total stellar mass within the NSC and NSD to be $5.3 \\times 10^9$~$M_{\\solar}$, which is unacceptably large. We hence ignore the latter possibility, and discuss based on the former possibility.  Origin of Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX therefore cannot be the same as that in the GRXE. \\citet{Ya09} found that the line intensity ratio of Fe \\emissiontype{XXVI} K$\\alpha$ to Fe \\emissiontype{XXV} K$\\alpha$ in $|l_*| < \\timeform{0D.3}$ is higher than that in $|l_*| > \\timeform{8D}$, which also supports  that origin of Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX may be different from that in the GRXE. \\citet{Re09} resolved more than 80\\% of the Fe K$\\alpha$ flux at $(l,\\ b)=(\\timeform{0D.1}, \\timeform{-1D.4})$ into point sources. Although this fact supports that origin of Fe \\emissiontype{XXV} K$\\alpha$ line in the GRXE is due to point sources, that in the GCDX is not yet answered. As is shown by the arrows in figures 4b and 4d,  Fe \\emissiontype{XXV}~K$\\alpha$ from this region is due to the sum of nearly equal contribution of the GD and GB (the GRXE components) and not the GCDX components (the NSC and NSD).  If we assume that origin of Fe \\emissiontype{XXV} K$\\alpha$ line in the GCDX is an ensemble of point sources with the same populations in the GRXE region, the iron abundances in the GCDX should be $3.8(\\pm0.3$)--$19(\\pm6$) times larger than that in the GRXE. This is unlikely because the abundance of iron has been reported at most two times of solar in the GCDX region (e.g. \\cite{Ko07b,No10}).  Another possibility is that the point source population in the GCDX is different from that in the GRXE. We see the Fe \\emissiontype{XXV} K$\\alpha$ intensity excess in the positive $l_*$ side (east) compared to the negative side (west) at $|l_*|=\\timeform{0D.1} $--$\\timeform{0D.3}$ (see figure 3a). This excess is globally associated to the clusters of high mass stars such as the Arches and Quintuplet clusters. \\citet{Mu09} also reported that the number density of X-ray stars in these regions exceeds that in the surroundings. They proposed that the excess is due to massive stars associated to the high activity of star formation in this region. The most probable massive stars are WR or OB stars. These stars often show moderately large equivalent width of Fe \\emissiontype{XXV} K$\\alpha$ ($EW_{6.7}$) similar to the GCDX (e.g. \\cite{Sk07, Hy08}). These massive stars might contribute to not only the Arches and Quintuplet clusters but also the entire GCDX. These stars are relatively bright with the typical X-ray luminosity of  $\\sim 10^{34}$ erg s$^{-1}$ \\citep{Po02, Wa06}. The deep Chandra observation revealed that point sources with the luminosity larger than $10^{32}$ erg s$^{-1}$  contribute less than 10\\% of the GCDX (see figure 6 of \\cite{Re07GC}). We therefore concluded that massive stars do not mainly contribute to Fe \\emissiontype{XXV} K$\\alpha$  line in the GCDX.  On the other hand, it may be conceivable that faint X-ray point sources with the flux below the current detection limits largely contribute to the GCDX. The most probable candidate sources for the Fe \\emissiontype{XXV} K$\\alpha$ line are cataclysmic variables (CVs), because their spectra are represented by high temperature plasma with large $EW_{6.7}$ and are associated with the 6.4 and 7.0 keV lines, very similar to those of the GCDX in the Fe K$\\alpha$ line band. The mean $EW_{6.7}$ for X-ray bright CVs in the luminosity range of $10^{30-33}$ erg s$^{-1}$ is about 200 eV \\citep{Ez99, Ra06}, while those  of the GCDX and GRXE are 450--550 eV \\citep{Ya09}. Thus even for the GRXE, the faint (unresolved X-ray) CVs should have, at least, two times larger $EW_{6.7}$ than the bright CVs. In addition, the number density of such (X-ray) faint CVs relative to near infrared stars should be $3.8(\\pm0.3$)--$19(\\pm6$) times larger than that in the GRXE to explain the observed Fe \\emissiontype{XXV} line excess in the GCDX. This possibility does not violate the total mass constraint of the NSC and NSD ($\\sim1.4 \\times 10^9$ $M_{\\solar}$), because the mass contribution of CVs to that of the SMD (mass of infrared stars) is only a small fraction (see e.g., \\cite{Pr07} and references there in). For relatively bright CVs, however, such number density excess in the GCDX region has not been reported \\citep{Re08}. For faint CVs,  studies for the spatial profile have not been done well. In order to settle these potential problems in the point source origin, deep observations, with lower detection limit than the current value by orders of more than 1--2 are required.  An alternative possibility is that the majority of the GCDX (NSC and NSD) comes from diffuse optically thin thermal plasma as already proposed  by \\citet{Ko07b}. Although the total flux of the Fe \\emissiontype{XXV} K$\\alpha$ line in the NSC and NSD regions is hardly explained by point sources, the spatial distributions are similar to the shapes of the NSC and NSD. The diffuse plasma, therefore, should be closely related to the ensemble of point sources. One possibility is that the plasma is made by multiple supernova explosions. This may be conceivable because the NSC and NSD regions are distinguished from the other regions (GB and GD) in respect to supernova activities; i.e., very high density of  molecular gas and on-going star formation.   \\begin{figure*}[htbp] \\begin{center} \\begin{center} (a)\\FigureFile(72mm,54mm){figure4a.eps} (b)\\FigureFile(72mm,54mm){figure4b.eps} (c)\\FigureFile(72mm,54mm){figure4c.eps} (d)\\FigureFile(72mm,54mm){figure4d.eps} \\end{center} \\caption{(a) The Fe \\emissiontype{XXV} K$\\alpha$ line intensity distribution along the Galactic plane as a function of $l_*$ (distance from Sgr A$^*$). For simplicity, only negative $l_*$ region is shown, except for those at  $|l_*| > \\timeform{8D.0}$ because the model in these region (Galactic disk) is symmetric with respect to the positive and negative $l_*$.  The green, purple, orange and yellow lines are the best-fit SMD model of the nuclear stellar cluster (NSC), nuclear stellar disk (NSD), Galactic disk (GD) and Galactic bulge (GB), respectively.  The black solid line is the sum of the four components of NSC, NSD, GD and GB. The emissivities of these components are fixed to be the same. The data of $|l_*|>\\timeform{1D.5}$ and $|b_*|>\\timeform{0D.5}$ regions are used simultaneously in the fitting.\\newline (b) The same as the left panel (a), but those along the Galactic longitude at $l_*=\\timeform{-0D.114}$ as a function of $b_*$ (distance from the Galactic plane). The arrow shows the position of $b=\\timeform{-1D.4} $ which is almost the same region as \\citet{Re09}. For simplicity, only negative $b_*$ region is  shown.\\newline (c) The same as the upper panel (a) but with the best-fit model in the whole region (GRXE and GCDX) with free parameters of the emissivities (or stellar mass densities) of NSC and NSD.\\newline (d) The same as the upper panel (b) but with the best-fit model same as the panel (c).\\newline } \\label{fig:SMM_fit} \\end{center} \\end{figure*}  \\begin{table} \\caption{Best-fit result with the stellar mass distribution model.} \\label{tab:SMM_fit} \\begin{center} \\begin{tabular}{cc} \\hline \\hline Component& Mass emissivity of \\\\ &the Fe \\emissiontype{XXV} K$\\alpha$ line $\\varepsilon$  ${}^*$\\\\ \\hline Nuclear stellar cluster & 111 $\\pm$ 37\\\\ Nuclear stellar disk  & 22 $\\pm$ 2 \\\\ Galactic disk$^\\dagger$ & 5.85 (fixed) \\\\ \\hline \\end{tabular} \\end{center} ${}^*$ See equation \\ref{eq:integral2}. The units are $10^{33}$ photons s$^{-1}$ $M_{\\solar}^{-1}$. ${}^\\dagger$ Mass emissivities of the Galactic bulge and the Galactic disk are fixed to the best-fit of $|l_*|>\\timeform{1D.5}$ and $|b_*|>\\timeform{0D.5}$ regions. \\end{table}"
    },
    "1107/1107.0516_arXiv.txt": {
        "abstract": "Using a nonparametric function estimation methodology, we present a comparative analysis of the WMAP 1-, 3-, 5-, and 7-year data releases for the CMB angular power spectrum with respect to the following key questions: (a) How well is the  power spectrum determined by the data alone? (b) How well is the $\\Lambda$CDM model supported by a model-independent, data-driven analysis? (c) What are the realistic uncertainties on peak/dip locations and heights? Our results show that the height of the power spectrum is well determined by data alone for multipole $l$ approximately less than 546 (1-year), 667 (3-year), 804 (5-year), and 842 (7-year data). We show that parametric fits based on the $\\Lambda$CDM model are remarkably close to our nonparametric fits in $l$-regions where data are sufficiently precise. In contrast, the power spectrum for an H$\\Lambda$CDM model gets progressively pushed away from our nonparametric fit as data quality improves with successive data realizations, suggesting incompatibility of this particular cosmological model with respect to the WMAP data sets. We present uncertainties on peak/dip locations and heights at the 95\\% ($2 \\sigma$) level of confidence, and show how these uncertainties translate into hyperbolic ``bands\" on the acoustic scale ($l_A$) and peak shift ($\\phi_m$) parameters. Based on the confidence set for the 7-year data, we argue that the low-$l$ up-turn in the CMB  power spectrum cannot be ruled out at any confidence level in excess of about 10\\% ($\\approx 0.12 \\sigma$). Additional outcomes of this work are a numerical formulation for minimization of a noise-weighted risk function subject to monotonicity constraints, a prescription for obtaining nonparametric fits that are closer to cosmological expectations on smoothness, and a method for sampling cosmologically meaningful power spectrum variations from the confidence set of a nonparametric fit. ",
        "introduction": "\\label{intro}  The angular power spectrum of cosmic microwave background (CMB) temperature fluctuations is a measurable physical quantity that depends sensitively on the physics of the early universe. In particular, the shape of the angular power spectrum and the locations and heights of its peaks relate directly to parameters in the underlying cosmological models. As such, it has been used extensively as an acid test of the relative merit of competing cosmological models, and as a rich source of information about cosmological parameters themselves.  Traditionally, and almost exclusively, cosmologists have resorted to model-based parametric statistical methods for estimating the CMB angular power spectrum from data. Parametric regression methods require the functional form of the unknown regression function $f$ to be pre-specified. The adjustable parameters in $f$, finite in number that is independent of the data size, are usually estimated by maximizing an appropriate likelihood function or a posterior distribution. In the cosmological context, it is conventional to employ parametric models that attempt to capture the essential physics of the problem via the pre-specified functional form, and any pre-existing knowledge about parameters is incorporated in the estimation process via appropriate prior distributions.  Nonparametric function estimation methods, on the other hand, assume no specific functional form for $f$, except for mild regularity conditions such as smoothness assumptions, membership to a function space, etc. In this approach, the number of parameters that define the unknown regression function $f$ is either infinite or grows proportionally with the data size, and the estimate $\\widehat{f}_N$ of $f$ is obtained by balancing bias and variance of $\\widehat{f}_N$ via optimal smoothing. Nonparametric methods are therefore model-independent, and are based on fewer and less restrictive assumptions about $f$. This, in turn, implies that any inferences about $f$ made from nonparametric regression analyses tend to be more data-driven as opposed to being primarily model-driven. In other words, to a greater extent nonparametric function estimation methods tend to infer what \\emph{is} rather than what \\emph{should be}. As such, nonparametric regression methods can be meaningfully employed as sanity-enforcing mechanisms on parametric analyses, thereby making the conclusions drawn more conservative. For example, a feature seen in a parametric fit that survives in a nonparametric analysis is likely to be a real and robust feature of the data itself, and not merely an artifact that is seen because a parametric model expects it to be there.  Alternative methodologies, such as the nonparametric methodology \\citep{GMN+2004,BSM+2007} used in this work, may allow posing inferential questions that are difficult to address using conventional methods. For example, this particular nonparametric methodology allows validating model-based, parametric fits against the confidence set for the nonparametric fit to the same data, possibly to rule them out as candidates for the true but unknown regression function. This methodology also has the formal advantage of being able to provide realistic uncertainties on \\emph{any} number of features of the angular power spectrum that are simultaneously valid at the same level of confidence. Such desirable features are arguably lacking in most mainstream methodologies, including Bayesian, that are commonly used in cosmology. (An incisive, insightful, and discerning discussion about the relative merits of this methodology over statistical methods conventionally employed in cosmology can be found in the two references cited above, and is best read in the original.)  The four CMB angular power spectrum data sets \\citep{HSV+2003,HNB+2007,NDH+2009,LDH+2011} released by the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{BBH+2003} mission, representing cumulative observations at the end of 1, 3, 5, and 7 years of operation, present a unique opportunity for statistical analysis. For example, till date these are claimed to represent the most precise and extensive full-sky CMB measurements ever made (only to be superseded by the Planck mission \\citep{TMP+2010}). The four data sets represent the evolution of data over a period of about a decade, thereby making it possible to assess progressive and possibly systematic resolution of features of the spectrum. From a statistical perspective, each of these moderately large data sets (minimum of 899 data points for the WMAP 1-year release) is not only heteroskedastic, but also has substantial correlations that arise in typical data pipelines \\citep{Tegmark1997,HBB+2003,HWH+2009,JBG+2007,JBD+2011}.  In this paper, we present a comparative nonparametric analysis of the WMAP 1-, 3-, 5-, and 7-year data sets for the angular power spectrum. Specifically, for each data realization, we address the following three key questions: (a) How well is the angular power spectrum determined by the data alone? (b) How well is the $\\Lambda$CDM model supported by a model-independent, nonparametric, data-driven analysis? (c) What are the realistic uncertainties on peak/dip locations and heights? Our analysis is based on a nonparametric function estimation methodology \\citep{GMN+2004,BSM+2007}, which is discussed in Sec.\\ \\ref{method} together with our extensions; i.e., a numerical formulation for minimization of a inverse-noise-weighted risk function subject to monotonicity constraints, a prescription for obtaining nonparametric fits that are closer to cosmological expectations on smoothness, and a method for sampling cosmologically meaningful power spectrum variations from the confidence set of a nonparametric fit. Results are presented in Sec.\\ \\ref{results}, and conclusion in Sec.\\ \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper, we have presented a comparative nonparametric analysis of the WMAP 1-, 3-, 5-, and 7-year data releases for the CMB angular power spectrum, using a nonparametric function estimation methodology \\citep{GMN+2004,BSM+2007}. In the context of this methodology, we have also presented our own numerical formulation for minimization of the inverse-noise-weighted risk function subject to monotonicity constraints, and a prescription for obtaining monotone nonparametric fits that are closer to cosmological expectations on smoothness. For all data realizations, we have presented results pertaining to the following questions: (a) how well is the angular power spectrum determined by the data alone, (b) how well is the $\\Lambda$CDM model supported by a model-independent, nonparametric, data-driven analysis, and (c) what are the realistic uncertainties on peak/dip locations and heights.  The motivation for the analysis presented here was to explore what could be inferred about the CMB angular power spectrum in a model-independent, data-driven manner. On the other hand, the basic physics of the CMB is quite well established. It would therefore be useful to connect a nonparametric/model-independent analysis such as ours with the known physics of the CMB angular power spectrum. This is reserved for the future.  To conclude, we have demonstrated in this paper the threefold utility of the nonparametric methodology used here for cosmological function estimation problems: as a method with sound formal guarantees, as a sanity-enforcing mechanism on parametric model-based analyses, and as a method that allows interesting inferential questions to be addressed and answered in a data-driven manner."
    },
    "1107/1107.5043_arXiv.txt": {
        "abstract": "We present the discovery and follow-up observations of SN~2008jb, a core-collapse supernova in the southern dwarf irregular galaxy ESO~302$-$14 \\mbox{($M_{B}=-15.3$~mag)} at \\mbox{$9.6$~Mpc}. This nearby transient was missed by galaxy-targeted surveys and was only found in archival optical images obtained by the Catalina Real-time Transient Survey and the All-Sky Automated Survey. The well sampled archival photometry shows that SN~2008jb was detected shortly after explosion and reached a bright optical maximum, $V_{\\rm max} \\simeq 13.6$~mag ($M_{V,\\rm max}\\simeq -16.5$). The shape of the light curve shows a plateau of $\\sim 100$~days, followed by a drop of $\\sim 1.4$~mag in $V$-band to a slow decline with the approximate $^{56}$Co decay slope.  The late-time light curve is consistent with $0.04\\pm 0.01$~M$_{\\odot}$ of $^{56}$Ni synthesized in the explosion. A spectrum of the supernova obtained 2~years after explosion shows a broad, boxy H$\\alpha$ emission line, which is unusual for normal type~IIP supernovae at late times. We detect the supernova in archival Spitzer and WISE images obtained $8-14$~months after explosion, which show clear signs of warm ($600-700$~K) dust emission. The dwarf irregular host galaxy, ESO~302$-$14, has a low gas-phase oxygen abundance, $\\rm 12 + log(O/H) = 8.2$ ($\\sim 1/5$~Z$_{\\odot}$), similar to those of the SMC and the hosts of long gamma-ray bursts and luminous core-collapse supernovae. This metallicity is one of the lowest among local ($\\lesssim 10$~Mpc) supernova hosts. We study the host environment using GALEX far-UV, $R$-band, and H$\\alpha$ images and find that the supernova occurred in a large star-formation complex. The morphology of the H$\\alpha$ emission appears as a large shell ($R \\simeq 350$~pc) surrounding the FUV and optical emission. Using the H$\\alpha$-to-FUV ratio, and FUV and $R$-band luminosities, we estimate an age of $\\sim 9$~Myr and a total mass of $\\sim 2\\times 10^{5}$~M$_{\\odot}$ for the star-formation complex, assuming a single-age starburst. These properties are consistent with the expanding H$\\alpha$ supershells observed in many well-studied nearby dwarf galaxies, which are tell-tale signs of feedback from the cumulative effect of massive star winds and supernovae. The age estimated for the star-forming region where SN~2008jb exploded suggests a relatively high-mass progenitor star with initial mass $\\rm M \\sim 20$~M$_{\\odot}$, and warrants further study. We discuss the implications of these findings in the study of core-collapse supernova progenitors. ",
        "introduction": "Core-collapse supernovae are energetic explosions that mark the death of stars more massive than $\\approx 8-10$~M$_{\\odot}$. They are extremely important in several areas of astrophysics, including the nucleosynthesis of chemical elements, energy feedback that affects the evolution of galaxies, the formation of compact object remnants, the production of high-energy particles, as tracers of recent star-formation in galaxies, and as test cases of massive stellar evolution.  Supernovae discovered in nearby galaxies have been particularly important for testing our physical understanding of the explosions and to establish direct links between the events and their progenitors, which puts tight constraints on stellar evolution theory. The two best-studied nearby core-collapse supernovae, the peculiar type~II SN~1987A in the LMC and the type~IIb SN~1993J in M81, had progenitors that were detected in pre-explosion images. Both were very interesting cases that challenged theoretical expectations based on stellar evolution models that predicted red supergiants as their progenitor stars: SN~1987A had a $\\simeq 20$~M$_{\\odot}$ blue supergiant progenitor star (e.g., Arnett et al. 1989) and SN~1993J had a $\\simeq 15$~M$_{\\odot}$ yellow supergiant progenitor (e.g., Aldering et al. 1994), which turned out to be a massive binary system (e.g., Maund et al. 2004). After more than a decade of work, Smartt et al. (2009) presented a thorough study of nearby ($\\lesssim 30$~Mpc) type~II-Plateau supernovae with deep pre-explosion imaging (see also Li et al. 2005), which allowed them to constrain the progenitors of type~IIP supernovae to be red supergiants with initial main-sequence masses in the range $8-17$~M$_{\\odot}$. This result is in conflict with well-studied red supergiant samples in the Galaxy and the Magellanic Clouds (e.g., Levesque et al. 2006), which also contain massive supergiants with $\\rm M\\simeq 20-30$~M$_{\\odot}$, and a Salpeter Initial Mass Function (IMF) for high-mass stars.  This ``red supergiant problem\" can be alleviated if high-mass red or yellow supergiants produce other spectroscopic or photometric type~II supernova sub-types (e.g., Smith et al. 2009; Elias-Rosa et al. 2009, 2010), or if high-mass supergiants experience strong winds due to pulsations (e.g., Yoon \\& Cantiello 2010). Very recently, Walmswell \\& Eldridge (2011) have proposed circumstellar dust as a possible solution. It can also be explained by failed supernovae, an hypothesis that would have several deep implications (e.g., Kochanek et al. 2008). In addition, there are known selection biases that have not been discussed in depth or accounted for. Of the 18 nearby type~II supernovae with fairly secure progenitor (either single, binaries, or compact clusters) detections\\footnote{These are: SN~1987A, 1993J; SN~1961V from Kochanek et al. (2011) and Smith et al. (2011c); SN~1999ev, 2003gd, 2004A, 2004am, 2004dj, 2004et, 2005cs, and 2008bk from the sample of Smartt et al. (2009); SN~2005gl, from Gal-Yam et al. (2007) and Gal-Yam \\& Leonard (2009); SN~2008ax from Crockett et al. (2009); SN~2008cn and 2009kr from Elias-Rosa et al.  (2009, 2010); SN~2009md from Fraser et al. (2011); SN~2010jl from Smith et al. (2011b); and SN~2011hd from Maund et al. (2011), Szczygiel et al. (2011) and Van Dyk et al. (2011). We do not include here low-luminosity transients with type~IIn-like spectral features and progenitor detections in the mid infrared, like SN~2008S and NGC~300-OT (e.g., Prieto et al. 2008b).}, two of the 18 events were initially discovered by the galaxy-targeted Lick Observatory Supernova Search (LOSS; Filippenko et al. 2001), and 14 of the 18 were initially discovered by dedicated amateur astronomers that mostly concentrate their search efforts in big galaxies. Only 2 of these supernovae were found in dwarf galaxies (SN~1987A and 2008ax) and all the type~IIPs (10 of the 18) were found in the disks of spiral galaxies, mostly in grand-design spirals.  The existing selection bias against finding nearby core-collapse supernovae in dwarf galaxies could be important because $20-40\\%$ of the local star formation rate density is in galaxies with absolute magnitude $M_B > -18$~mag (e.g., James et al. 2008; Young et al. 2008; Williams et al. 2011).  Since metallicity is a key parameter in massive stellar evolution and stellar death (e.g., Prieto et al. 2008a; Modjaz 2011, and references therein), changes in the fractions of different type~II spectroscopic subtypes with metallicity (e.g., Arcavi et al. 2010) could affect our census of progenitors in local samples. Also, possible variations of the stellar IMF with environment (e.g., Meurer et al. 2009; but see, e.g., Myers et al. 2011 and references therein) can introduce complications to the progenitor analyses. Recently, Horiuchi et al. (2011) discussed in detail related effects of incompleteness in the existing supernova rate measurements and how they affect their use as star formation rate indicators.  In this paper we present the discovery, follow-up observations, and analysis of the nearby type~II SN~2008jb ($\\alpha = 03^{\\rm h}51^{\\rm m}44\\fs 66$, $\\delta = -38^{\\circ} 27\\arcmin 00\\farcs1$) in a southern dwarf-irregular galaxy at $\\sim 10$~Mpc (Prieto et al. 2011). This bright $13.6$~mag supernova was missed by targeted southern supernova surveys, like CHASE (Pignata et al. 2009) and amateur searches, mainly because the host galaxy ESO~302$-$14 is a fairly low-luminosity dwarf and was not included in the catalogs of galaxies that are surveyed for supernovae. We are able to recover the optical light curves of this supernova from two surveys that do not target individual galaxies, but rather scan large areas of the sky: the Catalina Real-time Transient Survey\\footnote{{\\tt http://crts.caltech.edu}} (CRTS; Drake et al. 2009; Djorgovski et al. 2011) and the All-Sky Automated Survey (ASAS; Pojmanski 1998). In Section~\\S\\ref{sec1} we discuss the observations, including the optical discovery and follow-up, archival mid-infrared imaging, optical spectroscopy, and other archival data. In Section~\\S\\ref{sec2} we present detailed analysis and discussion of the optical and mid-IR observations. In Section~\\S\\ref{sec3} we discuss the results and conclusions. We adopt a distance of 9.6~Mpc to ESO~302$-$14 (Lee et al. 2009), estimated from the Virgo-inflow corrected recession velocity and $H_0=75$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sec3}  We have presented the discovery, follow-up observations, and analysis of SN~2008jb, a bright type~II supernova in the metal-poor, southern dwarf irregular galaxy ESO~302$-$14 at $\\sim 10$~Mpc. This $13.6$~mag supernova was found in archival data obtained by the CRTS and ASAS all-sky surveys. This transient was missed by galaxy-targeted supernova surveys like CHASE and by amateur astronomers mainly because the host galaxy is a low-luminosity dwarf, $\\sim 1$~mag fainter than the SMC, and targeted surveys use catalogs that are incomplete for small galaxies (e.g., Leaman et al. 2011).  SN~2008jb has $V$- and $I$-band light curves similar to normal type~IIP supernovae, with peak magnitude $M_V \\simeq -16.5$~mag, but it can also be classified as an intermediate case between type~IIP and type~IIL due to its faster initial decay, perhaps similar to SN~1999ga (Pastorello et al. 2009b). It shows a $\\sim 95$~day plateau and a fast $\\sim 1.4$~mag decline to a late-time decline slope of $0.013$~mag/day in the $V$-band. This decline is consistent with the radioactive decay of $^{56}$Co to $^{56}$Fe, and argues for $0.04 \\pm 0.01$~M$_{\\odot}$ of $^{56}$Ni synthesized in the explosion that is powering the lightcurve. We detect mid-IR emission from SN~2008jb $8-14$ months after explosion in three epochs of archival Spitzer and WISE data. The mid-IR emission has an SED with black-body temperature of $600-700$~K, characteristic of the warm dust emission seen in some nearby core-collapse supernovae and luminous transients (e.g., Kotak et al. 2009; Prieto et al. 2009). The evolution of the mid-IR emission with time is consistent with the products of radioactive decay heating the dust. The characteristic mid-IR dust radius shrinks with time, an evolution that is not typically seen in normal type~IIs (some exceptions are SN~2007it and SN~2007od; Andrews et al. 2010, 2011).  We obtained a spectrum of SN~2008jb about 2 years after the explosion. It displays a very broad (${\\rm FWHM}\\simeq 14000$~km~s$^{-1}$), boxy, and flat-topped H$\\alpha$ emission line, leading to its type~II supernova spectroscopic classification. The broad and boxy line profile seen in H$\\alpha$ is quite unusual for normal type~IIP supernovae at late times, but has been seen in some objects like SN~1993J, SN~2007od, and also a few well-studied type~IIL supernovae. We find that the H$\\alpha$ line luminosity is in excess of the expected luminosity from radioactive $^{56}$Co decay predicted by the models of Chugai (1991) for total ejected masses $\\rm M=5, 14, 20$~M$_{\\odot}$. This indicates that there is an external source of energy, like ejecta-CSM interaction, and/or the mass of the progenitor star is $\\gtrsim 20$~M$_{\\odot}$. We do not see clear signs of ejecta-CSM interaction like narrow lines or irregularities in the H$\\alpha$ profile. It would be interesting to obtain late time X-ray and radio observations in order to have independent constraints on the importance of ejecta-CSM interaction in SN~2008jb.  We studied the host galaxy environment of SN~2008jb in ESO~302$-$14 with optical spectra and imaging. Using the spectrum of an \\ion{H}{2} region at $\\sim 150$~pc from the supernova site we measure an oxygen abundance of $\\rm 12+log(O/H) = 8.21 \\pm 0.03$ (PP04~$O3N2$ method) and $\\rm 12+log(O/H) = 8.13 \\pm 0.03$ (PP04~$N2$ method) from the strong nebular emission lines, which is similar to the SMC and one of the lowest measured metallicities of local core-collapse supernovae environments (e.g., in the lower 3\\% of measured oxygen abundances of the core-collapse sample of Anderson et al. 2010). The supernova exploded in a large star-forming complex with strong optical and GALEX FUV emission which is surrounded by a large H$\\alpha$ ring with $\\rm R \\simeq 350$~pc. The H$\\alpha$-to-FUV ratio of this region is consistent with a stellar population with an age of $\\sim 9$~Myr derived from Starburst99 modeling and single-age stellar population models. This age implies a supernova progenitor mass of $\\simeq 23$~M$_{\\odot}$, assuming a single star (but see, e.g., Smith et al. 2011a for the importance of binary progenitors), if the progenitor formed in this star-forming episode. The equivalent width of the H$\\alpha$ and H$\\beta$ emission lines in the spectrum of the \\ion{H}{2} region give another constraint on the age of the region of $5-6$~Myr, $3-4$~Myr younger than the estimate from H$\\alpha$-to-FUV ratio. We note that the star-formation history could be more complicated than a single burst, and we plan to study the region in detail using multiwavelength data in a future study. In particular, it would be interesting to include high-resolution data from HST.  Large expanding H$\\alpha$ shells (supershells with radii $>300$~pc) have been observed and studied in many nearby star-forming dwarf galaxies and have typical dynamical ages of $\\sim 10$~Myr (e.g., Martin 1998), which is fairly consistent with the age we derive here from the H$\\alpha$ and FUV emission. These structures are produced by the combined effect of many supernova explosions and winds from massive stars (e.g., Chakraborti \\& Ray 2011), and are the likely precursors of galactic winds. In a sense, we may perhaps be witnessing supernova feedback in real-time in this star-forming region.  The bias to large star-forming galaxies is clearly present in the samples of nearby \\mbox{($\\lesssim 30$~Mpc)} core-collapse supernovae used for progenitor studies (e.g., Smartt et al. 2009), and has been discussed in detail as a possible explanation for the discrepancy between measured local supernova rates and predicted supernova rates from galaxy star-formation rates (e.g., Horiuchi et al. 2011). Since the environments of nearby core-collapse supernovae used for progenitor studies generally miss dwarf galaxies because they were not included in the original searches, the progenitor properties and conclusions drawn from these samples regarding stellar evolution are not complete. In particular, the dearth of high-mass ($\\sim 20-30$~M$_{\\odot}$) progenitor stars of type~IIP supernovae (``red supergiant problem\") could be alleviated if these progenitor stars prefer lower-metallicity environments. For example, this could be caused by evironmental variations in the stellar IMF (e.g., Meurer et al. 2009) or by changes in the fraction of type~II spectroscopic subtypes as a function of metallicity due to stellar evolution (e.g., Arcavi et al. 2010). Indeed, we find that the properties of the spectrum of SN~2008jb are more consistent with a massive progenitor (but see discussion about supernova modeling in, e.g., Smartt et al. 2009; Bersten \\& Hamuy 2009; Bersten et al. 2011), and the star-forming region where it was found has a young age compared with the ages of detected type~IIP progenitors ($\\gtrsim 15$~Myr).  Interestingly, a strong preference for low-metallicity hosts is observed in long GRBs (e.g., Stanek et al. 2006) and luminous core-collapse supernovae (e.g., Neill et al. 2011; Stoll et al. 2011), which have been linked with massive star progenitors ($\\rm M \\gtrsim 20-30$~M$_{\\odot}$). SN~2008jb offers the unique chance of studying in detail a nearby type~II supernova with host properties similar to long GRBs and the most luminous core-collapse supernovae.  The mapping between different classes of massive stars and their supernovae is not yet fully understood. Special insights are expected to be obtained when unusual explosions can be connected to unusual progenitor stars and galaxy hosts. Nearby objects are especially useful in terms of larger fluxes for an extended time after explosion, better spatial resolution for progenitor studies, and improved prospects for detection by new messengers like gamma rays (e.g., Timmes \\& Woosley 1997; Horiuchi \\& Beacom 2010), neutrinos (e.g., Ando \\& Beacom 2005; Kistler et al. 2011), and gravitational waves (e.g., Ott 2009). In addition, data from these objects are needed for a comprehensive understanding of the nearby universe.  It is difficult to find the nearest supernovae in small host galaxies with searches that target individual (generally large) galaxies, like LOSS, CHASE, and also amateur efforts. And it is difficult to find them with volume-based searches such as PTF (Rau et al. 2009) and Pan-STARRS (Kaiser et al. 2002) that have a deep, but relatively small survey area with good cadence. A shallower all-sky survey with excellent cadence, like ASAS, will help us find nearby \\mbox{($\\lesssim 30$~Mpc)} supernovae in all kinds of environments, including low-metallicity dwarf galaxies like the host of SN~2008jb (see also Khan et al. 2011, Stoll et al. 2011, for other supernovae studied with ASAS). Upgrades that will significantly increase the sensitivity and response speed of ASAS are underway."
    },
    "1107/1107.0982_arXiv.txt": {
        "abstract": "We present evidence for ultraviolet/optical microlensing in the gravitationally lensed quasar Q\\,0957+561.  We combine new measurements from our optical monitoring campaign at the United States Naval Observatory, Flagstaff (USNO) with measurements from the literature and find that the time-delay-corrected $r$ band flux ratio $m_{A} - m_{B}$ has increased by $\\sim 0.1$\\,magnitudes over a period of five years beginning in the fall of 2005.  We apply our Monte Carlo microlensing analysis procedure to the composite light curves, obtaining a measurement of the optical accretion disk size, $\\log \\{(r_{s}/\\textrm{cm})[\\cos (i)/0.5]^{1/2}\\} = 16.2 \\pm 0.5$, that is consistent with the quasar accretion disk size -- black hole mass relation. ",
        "introduction": "The history of microlensing in the $z_{s}=1.405$ quasar Q\\,0957+561 (hereafter Q0957), the first confirmed gravitational lens system, is paradoxically simple yet complex.  The possibility of microlensing in this double-image, wide-separation ($\\sim 6\\arcsec$) quasar was proposed by \\citet{chang79} soon after its discovery by \\citet{walsh79}.  Controversy over the time delay, currently accepted to be $\\Delta t_{AB} = 417\\,$days \\citep[e.g.,][]{vanderriest89,schild90,kundic97,shalyapin08}, spawned numerous time variability studies of Q0957 at optical and radio wavelengths which produced a wealth of monitoring data. Microlensing of the quasar on long time scales was soon established by a change in the time delay-corrected magnitude difference of the two images of $\\Delta (m_{A} - m_{B}) \\sim 0.25$ over the course of five years, observed by \\citet{vanderriest89} and confirmed by \\citet{pelt98}. However, after that initial event, occurring between 1982--1986, the difference light curve (the difference between the light curve for image A and the time-delay shifted light curve for image B) became nearly constant. Debate over the existence of measurable variability in the flux ratio percolated through the literature, with some authors reporting variability in the flux ratio on the order of several hundredths of a magnitude over timescales of weeks, days, and hours, while other groups simultaneously concluded that no microlensing with amplitude $ | \\Delta (m_{A} - m_{B}) |  \\gtrsim 0.05$ was detected \\citep[see, e.g.,][for summaries and references]{gilmerino01,schmidt10}. As the dispute surrounding the existence of short-timescale microlensing evolved, the absence of long-time scale microlensing became increasingly evident from the difference light curve's steadiness over roughly 20 years.  The only report of a possible long-time scale event after 1986, that of \\citet{ovaldsen03} of 5\\%-level variability over 300~days, was not corroborated independently by \\citet{wambsganss00} with data taken over the same time period.  Studies as recent as \\citet{shalyapin08} find no evidence for microlensing in their data.  With the development of numerical simulation methods to analyze microlensing observations \\citep[e.g.,][]{kochanek04,bate08,blackburne11}, quasar microlensing has meanwhile become a powerful tool with which to quantitatively study the properties of lens galaxies and the structure of quasars.  Although quasars and individual stars in lens galaxies cannot be resolved by conventional telescopes, by modeling the amplitudes of microlensing fluctuations we can measure the size of the continuum emission region, the masses and velocities of the stellar microlenses, and the stellar-to-dark mass fraction in the lens galaxy \\citep[e.g.,][]{kochanek04,pooley07,morgan08,mediavilla09,dai10}. Through analysis of microlensing variability at multiple wavelengths, the surface brightness profile of the quasar may be measured \\citep[e.g.,][]{poindexter08,eigenbrod08}. We are now moving into an era in which theoretical models of accretion disk structure can be tested with significant samples of lensed quasars, resulting in the observational confirmation of a relation between central black hole mass and accretion disk size \\citep{morgan10}, as well as the temperature profile of the \\citet{shakura73} thin accretion disk model \\citep[e.g.,][]{poindexter08,anguita08,mosquera11}.  Since the central black hole in Q0957 lies at the high end of the quasar black hole mass function \\citep[$1.0\\times 10^{9}\\,\\msun$;][]{assef11}, Q0957 provides a potentially interesting and important data point for testing models of accretion disk structure.  However, the available analyses of Q0957's only widely-acknowledged microlensing event from 1982--1986 pre-dated the use of the newer, highly quantitative microlensing analysis techniques, and the qualitative methods employed in the published analyses of the event produce either limits or constraints too broad to be informative for current tests of quasar structure and studies of the surface mass density of lens galaxies. For example, \\citet{pelt98} found that the 1982--1986 microlensing event was consistent with a quasar radius of ``roughly'' $3\\times 10^{15}\\,\\textrm{cm}$ and could be explained by microlens masses down to $10^{-5}\\,\\msun$ with only a small fraction ($<5\\%$) in solar mass objects.  \\citet{refsdal00} also analyzed the 1982--1986 microlensing event and subsequent 8-year quiescent period in Q0957 using three different simple lensing mass distribution models, obtaining rough constraints on the range of microlens masses of $10^{-6} < M/\\msun < 5$ and an upper limit on the size of the quasar's optical continuum region of $R < 10^{16}$\\,cm.  Attempts to analyze the negligibly low microlensing variability in Q0957's difference light curve with early numerical simulations \\citep{schmidt98,wambsganss00} could not constrain the quasar size.  Such studies could only rule out microlens masses for a given quasar size, and in the end obtained results contradictory to analyses of the 1982--1986 microlensing event.  Thus, the lack of unambiguous microlensing events in Q0957 in the relatively recent era of computationally-intensive microlensing analysis has been unfortunate, preventing improved determinations of the quasar's size and microlens mass properties.  However, as the saying goes, ``good things come to those who wait.'' In this paper, we combine three new seasons of optical photometric monitoring of Q0957 with previously-published light curves to demonstrate the return of long-time scale uncorrelated variability in the light curves of Q0957.  We take advantage of this new microlensing event to place the best constraints yet on the size of the optical accretion disk and the mean mass of the microlenses using the Bayesian Monte Carlo analysis technique of \\citet{kochanek04}.  In \\S\\ref{sec:obs_data} and \\S\\ref{sec:lens_model} we present our new monitoring observations of Q0957 and discuss the methods we use to model the uncorrelated variability in the light curves.  We present our results in \\S\\ref{sec:results}.  Throughout our discussion we assume a flat cosmology with $\\Omega_{\\textrm{M}}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_{0}=70\\,\\kms~\\mpc^{-1}$ \\citep{hinshaw09}. ",
        "conclusions": "\\label{sec:results}  In Figure~\\ref{fig:vel_dist} we show the probability distribution for the scaled effective source plane (Einstein) velocity $\\hat{v}_{e}$ resulting from the Bayesian analysis of our Monte Carlo light curve simulations, along with the probability density for the true source velocity, $P(v_{e})$, used to determine the mean microlens mass $\\langle M \\rangle$ and physical source size $r_{s}$. The $\\hat{v}_{e}$ distribution has a median of $1600\\,\\kms$ and is wide, with 68\\% confidence range of $600\\,\\kms < \\hat{v}_{e} < 3500\\,\\kms$.  The distribution is also visibly asymmetric, with a substantial low-velocity tail.  The width and asymmetry indicate the wide range of feasible solutions to the microlensing variability in Q0957 and reflect the ease with which the system's light curve was fit by the Monte Carlo code.  Because the Monte Carlo simulations produce large numbers of acceptable fits and the $\\chi^{2}$ per degree of freedom for the best fits tends to be quite low ($\\sim 0.4$) even with insignificant (small) assumed values for the systematic errors of the flux ratio and photometry, we conclude that the spread in the velocity distribution is a consequence of the low intrinsic variability of Q0957 coupled with the very low amplitude of the microlensing signal, making the light curve simple to reproduce.   An additional factor may be that the lack of observational constraints in the four-year gap in the compiled light curve permits considerable freedom in model behavior, as shown in Figure~\\ref{fig:ml_model_fits}.   We tested this hypothesis by conducting an experiment in which we forced the flux ratio in the gap to stay within $\\pm 0.05$\\,mag of its mean value in the dates bracketing the gap, $(m_{A} - m_{B}) = 0.04 \\pm 0.03\\,\\textrm{mag}$.  The experiment resulted in a minor narrowing of the $\\hat{v}_{e}$ distribution, but the change was not significant enough to justify any firm conclusions about the influence of the gap on our results.  Because the mean microlens mass is proportional to the inverse square of the scaled effective velocity ($\\hat{v}_{e} = v_{e}\\langle M/\\msun \\rangle^{-1/2}$), the broad range in permitted velocities causes the mean mass to be poorly constrained, which is apparent in the probability distribution for $\\langle M/\\msun \\rangle$ shown in the right panel of Figure~\\ref{fig:vel_dist}. The median of the $\\langle M/\\msun \\rangle$ distribution is $0.2\\,\\msun$, with a 68\\% confidence range of $0.02\\,\\msun < \\langle M \\rangle < 1.3\\,\\msun$. Our estimate of the mean microlens mass is appropriate for stars in an old elliptical galaxy, and falls well within the most probable range found by \\citet{refsdal00} of $10^{-6} < M/\\msun < 5$.   Our constraints do not permit a dominant population of planetary and sub-planetary mass ($< 10^{-3}\\,\\msun$), compact microlenses in the lens galaxy, consistent with the results of \\citet{schmidt98} and \\citet{wambsganss00}.  Our new constraints are also not consistent with the result of \\citet{pelt98} for the 1982--1986 event, in which a significant population of stellar mass microlenses is ruled out. However, as our microlensing analysis technique is significantly more sophisticated than that used in earlier studies and does not rely on the assumption of a single value for the source size or effective velocity to derive the microlens mass, we consider our new result to be much more robust.  While we have attempted to constrain the stellar-to-dark mass fraction $\\kappa_{\\ast}/\\kappa$ in the lens galaxy of the Q0957 system, the probability distribution resulting from our Bayesian analysis is uninformative, with no strong peaks or trends favoring any particular $\\kappa_{\\ast}/\\kappa$ value. We suspect that the primary reason for this failure is that there is insufficient uncorrelated variability in the light curves of Q0957 to constrain the dark matter fraction.  However, the addition of X-ray data to the microlensing analysis may provide stronger constraints, as it has in the case of PG\\,1115+080 \\citep{morgan08,pooley09} and RXJ\\,1131-1231 \\citep{dai10}.  In Figure~\\ref{fig:size_dist} we show the probability distribution for the physical source size $r_{s}$ of the quasar's accretion disk in the observed-frame $r$-band which results from our microlensing analysis.   The distribution as shown has not been corrected for inclination $i$; however, in the text that follows all the numerical quantities we discuss will be corrected by a factor of $(\\cos i)^{-1/2}$ assuming an inclination of $60\\degr$, the expectation value of a random distribution of inclinations. Note that the thin disk scale radius can be converted to a half-light radius using the relation $r_{1/2} = 2.44\\,r_{s}$.  The median of the microlensing thin-disk size distribution is $\\log (r_{s}/\\textrm{cm}) = 16.2 \\pm 0.5$, where the error bar represents the bounds of the 68\\% confidence interval. When converted to a half-light radius, which is comparable between different source models \\citep{mortonson05}, the source size we obtain from our analysis [$\\log (r_{1/2}/\\textrm{cm}) = 16.5 \\pm 0.5$] appears to be marginally consistent with the $R$-band half-light radii (converted from $\\sigma$ of a Gaussian disk profile) obtained by \\citet[][$10^{15.5}\\,\\textrm{cm}$]{pelt98} and \\citet[][$< 10^{16.1}\\,\\textrm{cm}$]{refsdal00} from the data set covering the 1982--1986 microlensing event and the following 8 years of stable flux ratio. We note that our source size result is substantially larger than those used by \\citet{schmidt98} and \\citet{wambsganss00} to constrain the microlens mass distribution [$\\log (r_{s}/\\textrm{cm}) < 14.2$]; yet our mean microlens mass is still consistent with their constraints. We also experimented with the use of a uniform prior on the microlens mass of $0.1 < \\langle M/\\msun \\rangle < 1.0$. Although the median of the resulting $r_{s}$ distribution is essentially the same as that found without imposing the mass prior, within the uncertainties, the distribution yielded by the prior contains a small but significant probability in solutions with small source sizes which are not physically reasonable.  We suspect that the occurrence of this ``shelf\" at the low end of the size distribution is an artifact of a combination of the paucity of microlensing variability in the observations and the mass prior itself.  By limiting $\\langle M \\rangle$ to a range about which its distribution (derived from the $v_{e}$ prior) is not symmetric, we give more statistical weight to solutions from one side of the distribution.  Since the Einstein velocity ($\\hat{v}_{e}$) scale implied by a given microlens mass scales as $\\hat{v}_{e} = v_{e} \\langle M/\\msun \\rangle^{-1/2}$, the solutions with very low Einstein velocities are given more weight relative to the highest velocity solutions when the prior is applied.  In the case of Q0957, the solutions with low Einstein velocity correspond preferentially to small source sizes; however, at very small sizes all solutions become equally likely since none of them pass near a caustic.  \\citet{morgan10} used the same microlensing analysis technique as we use here on light curves of a sample of 11 different gravitationally lensed quasars to show that the quasar accretion disk size at 2500\\,\\AA\\ inferred from microlensing ($r_{2500}$) is correlated with central black hole mass as $r_{2500} \\propto M_{\\textrm{BH}}^{0.80 \\pm 0.17}$.  Although the scaling of the observed correlation is consistent within the uncertainties with thin disk theory ($R \\propto M_{\\textrm{BH}}^{2/3}$), the correlation implies a quasar radiative efficiency $\\eta$ that is approximately an order of magnitude lower than is expected based on the local supermassive black hole mass and quasar luminosity functions.  We can use this $r_{2500}$--$M_{\\textrm{BH}}$ correlation as an independent check on our assertion that the uncorrelated variability we have observed in Q0957 is a result of microlensing. To scale the observed-frame $r$-band source size we find for Q0957 to a rest-frame wavelength of 2500\\,\\AA, we assume the $R\\propto \\lambda^{4/3}$ scaling of thin disk theory which is supported by several observational studies of quasar microlensing \\citep[e.g.,][]{anguita08,poindexter08,floyd09,mosquera11}. In doing so, we note that the effect of the wavelength scaling is minimal for Q0957 since the effective wavelength of the observed-frame $r$ filter corresponds to 2593\\,\\AA\\ at $z=1.41$. For the black hole mass we use the C\\,\\textsc{iv} emission-line estimate of \\citet{assef11}, which has an associated systematic uncertainty of 0.33 dex.   In Figure~\\ref{fig:mbh_vs_r2500}, we place Q0957 on the best-fit $r_{2500}$--$M_{\\textrm{BH}}$ relation found by \\citeauthor{morgan10}, along with the quasar data points used to derive the correlation.  Q0957 is consistent with the relation; the agreement supports our assertion that microlensing is the source of the quasar's long-timescale uncorrelated variability.  Similarly to \\citet{morgan10}, we also show for comparison in Figure~\\ref{fig:mbh_vs_r2500} quasar source sizes calculated by two additional methods.  First, we calculate the theoretical scale radius of a thin accretion disk at a rest wavelength of 2500\\,\\AA\\ from its central black hole mass (``theory size\") using Eq.\\ 2 of \\citeauthor{morgan10}, and assuming accretion efficiency $\\eta = 0.1$ and Eddington ratio $L/L_{E} = 1/3$ \\citep{kollmeier06}.  Second, we calculate the disk size under the thin disk model (a blackbody with a $T\\propto R^{-3/4}$ temperature profile) constrained by the observed magnification-corrected optical flux/luminosity from \\emph{HST} F814W measurements (``flux size\", see Eq.\\ 3 of \\citeauthor{morgan10}). When we compare the results of the different methods of size calculation for Q0957, we observe that the microlensing size is larger than the theory size, also noted for other lensed quasars by \\citet{morgan10} and \\citet{blackburne11}. \\citet{pooley07} arrive at a similar conclusion, although they use bolometric luminosity-based black hole masses to calculate their theory sizes which are thus more similar to our flux sizes. Furthermore, the microlensing size for Q0957 is larger than its flux size as well, which was also found by \\citeauthor{morgan10}\\ and \\citet{mediavilla11} for their respective quasar samples. However, the significance of the discrepancy we find for Q0957 is not as large as that typically found by \\citet{blackburne11} for their sample: while they rule out the black hole mass-based theoretical prediction by at least $3\\sigma$ in nearly every case, the flux size for Q0957 falls within the 68\\% confidence interval (essentially $1\\sigma$) of the microlensing size, and the theory size within the 80\\% confidence interval ($\\sim 1.3\\sigma$). While this may simply reflect the relatively large uncertainty in the microlensing size estimate for Q0957 compared to other microlensed quasars, it is also possible that the difference in significance is related to our different treatment of systematic errors from the analysis of Blackburne et al.  The origin of the discrepancies between the microlensing size, theory size, and flux size remain unclear. As \\citeauthor{morgan10}\\ and \\citeauthor{blackburne11}\\ point out, adjusting the mass accretion rate ($L/\\eta L_{E}$) of Q0957 by lowering $\\eta$ or using a higher-than-typical Eddington ratio may resolve the discrepancy of the theory size with the microlensing size, but such adjustments will not address the fact that the flux size is smaller than the microlensing size.  One possible explanation for the size discrepancy is that the observed $r$-band flux may be contaminated by (a) UV/optical light from the continuum source scattered by the broad-line region (BLR); or (b) higher energy continuum emission reprocessed by the BLR and re-emitted as UV/optical emission lines.  Contamination from line emission is an especially strong concern for our observations of Q0957 because the quasar's redshifted Mg\\,\\textsc{ii} line falls within the $r$ bandpass and could cause us to overestimate the microlensing size. Thus, it is important that we consider contamination scenarios in our analysis.  So, we have repeated our Monte Carlo light curve simulations assuming that different percentages of the observed $r$-band flux can be attributed to unmicrolensed contamination from the BLR. When we do so, we find that if we assume 10\\% contamination we still obtain a median physical accretion disk size in $r$-band of $\\log (r_{s, 10}/\\textrm{cm}) = 16.1 \\pm 0.5$. Even when we assume a 30\\% contribution to the observed flux from contamination, we obtain $\\log (r_{s, 30}/\\textrm{cm}) = 16.0^{+0.5}_{-0.4}$, insignificantly different from our original microlensing size. Thus, neither scattering nor line contamination appear to produce a significant effect on the microlensing size result, relative to our other sources of uncertainty.  Moreover, even if line contamination produced a significant decrease in the microlensing size, the flux size would be reduced along with the microlensing size, and the two would still not be reconciled.  Although we are not able to resolve the discrepancy between the different accretion disk size estimates and our microlensing size measurement, we are nonetheless very encouraged to observe a long-term microlensing event in a quasar at the high end of the black hole mass function. Moreover, Q0957 is relatively easy to monitor with ground-based observatories, since the quasar images and lens galaxy are relatively widely separated and the time delay is well-determined.  Future multi-wavelength optical monitoring as well as X-ray monitoring will be most informative to examine the temperature profile of the accretion disk and compare the X-ray and optical sizes."
    },
    "1107/1107.1984_arXiv.txt": {
        "abstract": "{Observations of star forming environments revealed that the abundances of some deuterated interstellar molecules are markedly larger than the cosmic D/H ratio of 10$^{-5}$. Possible reasons for this pointed to grain surface chemistry. However, organic molecules and water, which are both ice constituents, do not enjoy the same deuteration. For example, deuterated formaldehyde is  very abundant in comets and star forming regions, while deuterated water rarely is. In this article, we explain this selective deuteration by following the formation of ices (using the rate equation method) in translucent clouds, as well as their evolution as the cloud collapses to form a star. Ices start with the deposition of gas phase CO and O onto dust grains. While reaction of oxygen with atoms (H or D) or molecules (H$_2$) yields H$_2$O (HDO), CO only reacts with atoms (H and D) to form H$_2$CO (HDCO, D$_2$CO). As a result, the deuteration of formaldehyde is sensitive to the gas D/H ratio  as the cloud undergoes gravitational collapse, while the deuteration of water strongly depends on the dust temperature at the time of ice formation. These results reproduce well the deuterium fractionation of formaldehyde observed in comets and star forming regions and can explain the wide spread of deuterium fractionation of water observed in these environments.}\\rm ",
        "introduction": "Stars form through the collapse of interstellar clouds which are composed of gas and dust. As the cloud collapses, its density increases, its temperature decreases and it becomes shielded from external UV radiation that would otherwise dissociate chemical species present in the medium. In these regions, dust grains grow thick icy mantles from the deposition \\rm{of atomic and}\\rm\\ molecular species onto their surfaces. This has been confirmed by observations of starless cores (\\citealt{tafalla2006}), which show that most gas phase species suffer a significant drop towards the core center. The missing gas species constitute the icy mantles that cover dust grains. As a star forms and heats up its environment, \\rm{ a portion of the}\\rm\\ icy mantles are released into the gas phase. This phase of star formation, called the hot core (massive stars) or hot corino (low mass stars) phase, exhibits a very complex chemistry rich in oxygen and nitrogen bearing molecules (\\citealt{cazaux2003}), but also shows species such as H$_2$CO and methanol, which are building blocks for more complex organic molecules essential for the formation of life. Also, while deuterium is known to be 10$^{5}$ times less abundant than hydrogen (\\citealt{linsky2003}), many of these environments present  a chemistry rich in highly deuterated species. \\rm{In particular, deuterated methanol has been found in molecular clouds (\\citealt{parise2004}) and in low mass protostars, with a maximum of  $\\sim$ 30$\\%$ for CH$_2$DOH/CH$_3$OH  (\\citealt{parise2002}) and of $\\sim$ 3$\\%$ for CD$_3$OH/ CH$_3$OH (\\citealt{parise2004}), while no detection has been reported in comets (\\citealt{crovisier2004}). Formaldehyde present a high deuterium fractionation as well (see Table 1), HDCO/H$_2$CO is of a few percent towards protostellar cores (\\citealt{roberts2007}, \\citealt{bacmann2002}) and towards hot cores (\\citealt{turner1990}) and of a tenth of a percent towards hot corinos (\\citealt{parise2006}).}\\rm\\ Its doubly deuterated form D$_2$CO/H$_2$CO can reach up to 5$\\%$ in hot corinos (\\citealt{parise2006}). Species that have been observed with 3 deuterium atoms (ND$_3$; \\citealt{vanderTak2002}) and CD$_3$OH (\\citealt{parise2004}) are expected to be 10$^{15}$ times less abundant than their hydrogen form, \\rm{with respect to elemental abundances.}\\rm\\ Yet, their enhancement can be of 12 orders of magnitude (\\citealt{ceccarelli2007}).  This high degree of deuteration is attributed to grain surface chemistry. However,  while water co-exists with methanol and formaldehyde in ices, it rarely \\rm{has}\\rm\\ the same deuteration. In general, the degree of deuteration of organic material is very high compared to that measured in water (\\citealt{ehrenfreund2002}). The deuterium fractionation of water (Table 1) can be about  HDO/H$_2$O$\\sim$0.01$\\%$ towards massive hot cores (\\citealt{gensheimer1996}) and $\\sim$0.02-0.05$\\%$ in comets  and asteroids (\\citealt{altwegg2003}). These values are similar to SMOW (Standard Mean Ocean Water), which reports a fraction HDO/H$_2$O  of 0.015$\\%$ in our oceans. \\rm{However,  some detections showed a high HDO/H$_2$O ratio of few $\\%$ towards the hot corino NGC1333-IRAS2A (\\citealt{liu2011}), and IRAS16283-2422 (Coutens et al. in prep., \\citealt{parise2005}). }\\rm In this study, we address the selective deuteration of water compared to formaldehyde. We used H$_2$CO as a representative of organic molecules (H$_2$CO is a building block of most of the complex organic molecules found in space), but our study also apply to other species formed with CO, such as CH$_3$OH. First, we determine the chemical gas phase content of the region where ices form.  Then, we use a rate equation method that follows the accretion and reactions of gas phase species that yield ices. Finally, we simulate how the ice compositions change as they undergo gravitational collapse to form stars. These latest results are compared with observations of star forming environments as well as cometary material. ",
        "conclusions": "We used the rate equation method to simulate the formation of water and formaldehyde, as well as their isotopologues on dust grain surfaces. We model the formation of ices in translucent clouds and incorporate these ices in denser environments that undergo gravitational collapse to form stars.  We find that  HDCO/H$_2$CO and D$_2$CO/H$_2$CO depend on the gas D/H ratio which increases as the medium collapses and becomes denser.  For water, on the other hand, HDO/H$_2$O strongly depends on the dust temperatures at which ices form since water formation involves H$_2$ at low T$_{dust}$ (H$_2$ very abundant on dust surface) and H at higher T$_{dust}$ (H$_2$ evaporate). Consequently, the degree of water deuteration is directly related to the physical conditions (i.e., dust temperature) at which ices form in the quiescent cloud just before it collapses to form stars. \\rm{H$_2$CO is the simplest representative of organic species formed via hydrogenation of CO onto dust surfaces. Our results can then be generalized to other species (such as CH$_3$CO), which also depends on surface hydrogenation of CO. This explains the selective deuteration of water and organic molecules observed in the ISM, similar to that found in comets, IDPs and chondrites (\\citealt{ehrenfreund2002}, Fig. 9)\\rm.     \\subsection*"
    },
    "1107/1107.4045_arXiv.txt": {
        "abstract": "Peculiar velocities change the expansion rate of any observer moving relative to the smooth Hubble flow. As a result, observers in a galaxy like our Milky Way can experience accelerated expansion within a globally decelerating universe, even when the drift velocities are small. The effect is local, but the affected scales can be large enough to give the false impression that the whole cosmos has recently entered an accelerating phase. Generally, peculiar velocities are also associated with dipole-like anisotropies, triggered by the fact that they introduce a preferred spatial direction. This implies that observers experiencing locally accelerated expansion, as a result of their own drift motion, may also find that the acceleration is maximised in one direction and minimised in the opposite. We argue that, typically, such a dipole anisotropy should be relatively small and the axis should probably lie fairly close to the one seen in the spectrum of the Cosmic Microwave Background.\\\\\\\\ PACS numbers: 98.80.-k, 95.36.+x, 98.80.Jk ",
        "introduction": "\\label{sI} Realistic observers do not simply follow the smooth universal expansion, but have their own `peculiar' motions as well. The dipolar anisotropy of the Cosmic Microwave Background (CMB), in particular, has been interpreted as the result of our drift flow (at roughly 600~km/sec) relative to the cosmic rest-frame. Peculiar motions are believed to be the result of structure formation and are expected to fade away as we move on to progressively larger lengths. Thus, the current concordance (WMAP5-normalised) cosmological model predicts drift velocities of approximately 100~km/sec on scales larger than $50h^{-1}$~Mpc ($h$ is the Hubble parameter in units of 100~km/sec\\,Mpc). Recent observations however seem to indicate bulk flows with significantly larger amplitude and scale~\\cite{KA-BKE}-\\cite{LTMC}, casting some doubt on the $\\Lambda$CDM scenario~\\cite{P}. Here, we look into the potential implications of such large-scale drift motions, which are sometimes referred to as `dark flows', for the kinematics of the associated observers.  The first question is whether drifting observers in a perturbed, dust dominated Friedmann-Robertson-Walker (FRW) universe and those following the Hubble expansion could assign different values (and signs) to their respective deceleration parameters. Whether, in particular, it is theoretically possible for a peculiarly moving observer to `experience' accelerated expansion while the universe is actually decelerating. We find that the answer to this question is positive, when the peculiar velocity field adds to the Hubble expansion. In other words, the drifting observer should reside in a region that expands faster than the background universe. Then, around every typical observer in that patch, there can be a section where the deceleration parameter takes negative values and beyond which it becomes positive again. Moreover, even small (relative to the Hubble rate) peculiar velocities can lead to such local acceleration. The principle is fairly simple: two decelerated expansions (in our case the background and the peculiar) can combine to give an accelerating one, as long as the acceleration is `weak' (with $-1<q<0$ -- where $q$ is the deceleration parameter) and not `strong' (with $q<-1$) -- see \\S~\\ref{ssWSAE} below. Overall, accelerated expansion for a drifting observer does not necessarily imply the same for the universe itself. Peculiar motions can locally mimic the effects of dark energy. Furthermore, the affected scales can be large enough to give the false impression that the whole universe has recently entered an accelerating phase.  We then ask whether contemporary observations of large-scale bulk flows report drift velocities fast enough to support the above scenario. Again, the answer is positive. Although the $\\Lambda$CDM model predicts relatively weak peculiar motions, recent independent surveys seem to indicate larger rms bulk velocities on scales close to $100h^{-1}$~Mpc and beyond~\\cite{KA-BKE}-\\cite{LTMC}. The analysis of~\\cite{WFH}, for example, finds velocities greater than $400$~km/sec on the aforementioned lengths. Analogous results were also reached through the complementary methods of~\\cite{LTMC}, while faster bulk flows have been reported by~\\cite{KA-BKE}. The latter study, in particular, suggests almost constant peculiar velocities of roughly 1000~km/sec extending as far out as 800~Mpc, with no sign of converging beyond this threshold. On smaller lengths (between 30 and 60~Mpc), on the other hand, the work of~\\cite{LS} indicated a positive variance in the local Hubble rate of up to 10\\%. Using these measurements, we find that drifting observers can experience locally accelerated expansion on scales of several hundred Mpc (perhaps even larger), depending on the density of the region in question.  When all the observers in the drifting patch have peculiar velocities equal to the mean bulk flow of the region, they should see no preferred direction in the distribution of the deceleration parameter within the patch. In reality, however, observers have their own peculiar motions, which are generally different from the average. In principle, this should trigger a dipolar anisotropy in the measured distribution of the deceleration parameter, analogous to the CMB dipole. Now, however, the rest frame is not that of the universal expansion but the one defined by the mean flow of the patch. In other words, there should be a preferred axis in the accelerated domain, with the acceleration maximised in one direction and minimised in the opposite. Nevertheless, when dealing with typical observers -- those with peculiar velocities close to the average one, the aforementioned dipolar anisotropy must be small. Furthermore, the direction of the $q$-dipole should lie fairly close to that of the CMB, since they are both the result of the observers' drift flow. Interestingly, some recent reports seem to indicate that a `cosmological axis', namely a weak dipolar anisotropy aligned close to the CMB dipole, might actually reside within the supernovae data~\\cite{KL}.  We begin by outlining the basic features of the peculiar kinematics in linearly perturbed, dust-dominated FRW universes. Our first aim is to show that drifting observers measure different expansion rates (and therefore different deceleration parameters) from those following the background Hubble flow. We then explain how weak peculiar motions can lead to locally accelerated expansion within a globally decelerating universe. Section 3 provides the conditions for apparent universal acceleration and uses current observations to estimate the scale of the accelerated domain and the average value of the deceleration parameter there. In section 4 we consider anisotropic peculiar flows and obtain expressions for the main kinematic variables. These formulae are then used to discuss dipole-like patterns in the distribution of the deceleration parameter, which are caused by the peculiar motion itself. Finally, we close with a brief summary of our results and a general discussion. ",
        "conclusions": "\\label{sD} A little more than a decade ago, observations of high-redshift supernovae indicated that our universe was expanding at an accelerating pace~\\cite{Retal}. This conclusion was reached after applying the supernovae luminosity distances to the distance-redshift relation of an exact FRW model. The results have repeatedly returned negative values for $q$, indicating an accelerated expansion for our universe. In particular, the deceleration parameter was estimated close to -0.5. The same measurements also suggested that the accelerated phase was a relatively recent event, putting the transition from deceleration to acceleration around $z=0.5$~\\cite{TR}. As one looks at smaller redshifts, however, the picture seems to become less clear and there have been suggestions that the cosmic acceleration may have just peaked and started to slow down~\\cite{ST}. Explaining the supernovae results has dominated almost every aspect of contemporary cosmology. Dark energy, an unknown and elusive form of matter with negative gravitational energy, has so far been the most popular answer. There has been scepticism, however, since the dark-energy scenario lacks a satisfactory explanation in terms of fundamental physics and also faces an awkward coincidence problem (see~\\cite{S} for a discussion). Possible alternatives to dark energy, include modifying General Relativity, introducing extra dimensions, or abandoning the Friedmann models altogether.  The reason behind such drastic measures was that negative values for the deceleration parameter seemed impossible in conventional FRW cosmologies, as well as in perturbed FRW models. This is not necessarily the case however. Peculiar motions can locally induce weakly accelerated expansion (with $-1<\\tilde{q}<0$ -- see \\S~\\ref{ssWSAE}), even when the drift velocities are small and matter is simple pressure-free dust, namely within the limits of the linear approximation. Moreover, although the effect is local, the affected scales can be large enough (of the order of several hundred Mpc -- perhaps even larger) to give the false impression that the whole universe has recently moved into a phase of global acceleration.  Technically speaking, the physics discussed here stems from the fact that the Raychaudhuri equations in the two coordinate systems (the CMB and that of a drifting observer) are not the same. The difference is in the 4-acceleration term, which vanishes in the CMB frame but takes nonzero values in any other relatively moving reference system. In practice, this means that the expansion rate of the drifting observer and that of the background Hubble flow are generally different. It is then theoretically possible to experience local acceleration in a globally decelerating universe. Alternatively, one could say that the observer's total motion is the `sum' of the smooth Hubble expansion and of their own peculiar flow. Then, it is fairly easy to show that the two constituent motions (although both decelerating) can in principle combine to give weakly accelerated expansion (with $-1<\\tilde{q}<0$ -- see \\S~\\ref{ssWSAE}). The supernovae data seem to point towards weak acceleration.  There are, of course, certain requirements that need to be fulfilled, if peculiar motions are to cause local acceleration. The main one is that the drift flow should add to the Hubble expansion (i.e.~$\\vartheta=\\tilde{\\rm D}^av_a>0$). Put another way, the patch $\\mathcal{A}$ in Fig.~\\ref{fig:pvel} should be expanding faster (even slightly faster) than the universe itself. In that case, the two decelerating expansions (the background and the peculiar) can locally combine to give a weakly accelerating one. One might then ask what is the likelihood for the Milky Way to rest within a region like $\\mathcal{A}$. Given that the scales of interest are of cosmological relevance, the chances could be fairly high (perhaps higher than 50\\%). The reason is that systematically negative values of $\\vartheta$ are primarily associated with structures that have already decoupled from the background expansion, `turned around' and started to collapse. Such gravitationally bound systems, however, are believed to be considerably smaller in size.  The scale of the accelerated patch and the value of the deceleration parameter there, are decided by the speed of the drift flow and by the density of the matter in the section. In general, the larger the peculiar velocity and the lower the matter density, the larger the accelerated region and the lower (the more negative) the deceleration parameter. In our examples, we have adopted the peculiar velocities reported in the surveys of~\\cite{KA-BKE}. These suggest drift flows of approximately 1000~km/sec extending as far out as 800~Mpc (perhaps even further out). Based on these, we found that the accelerated patch could reach lengths of the order of 1000~Mpc. The deceleration parameter was found to drop with scale, with its value ranging between $-0.11$ and $-0.33$ close to 100~Mpc and tending towards zero as one moves on to progressively larger lengths. As far as the matter density is concerned, we have assumed values in the $(0.01,0.1)$ range for the effective density parameter of the accelerated patch. The use of underdense regions, which expand faster than the background universe, renders certain analogies between this scenario and the large-void models~\\cite{C}. Here, however, we do not need to reside near the centre of the void. Also, given that peculiar velocities are the result of structure formation, there are parallels with the `backreaction' scenarios as well (see~\\cite{R} for a recent account).  Drift flows are typically associated with dipolar anisotropies, triggered by the preferred spatial direction that the peculiar velocities introduce. The CMB dipole is believed to be an example of such a Doppler-like effect, caused by the motion of our Local Group relative to the rest-frame of the smooth Hubble expansion. Similarly, the motion of an observer in section $\\mathcal{A}$, relative to the mean bulk flow of that region, can induce a dipole-like anisotropy in the spatial distribution of the deceleration parameter. However, when the observer is a typical one -- with drift velocity close to the average flow of the patch -- the aforementioned dipolar anisotropy should be small. With these in mind we have considered anisotropic peculiar motions, analysed their kinematics and used the associated equations to discuss dipole-like anisotropies. Our results indicate that drift flows are generally consistent with some degree of dipolar anisotropy in the spatial distribution of $\\tilde{q}$. Moreover, if such a $\\tilde{q}$-dipole exists, the axis should not lie far from that of its CMB counterpart. This sounds plausible, given that both dipoles are induced by the drift motion of the Milky Way, though relative to different rest-frames in each case. Recently, there has been a number of reports suggesting that a small dipolar anisotropy might actually reside in the supernovae data~\\cite{KL}. Moreover, the axis of the $\\tilde{q}$-dipole appears to lie close to that of the CMB. In the literature, this has been largely treated as coincidental. Viewed from the perspective of this report, however, the alignment of the two dipolar axes seems more like a natural consequence rather than a mere coincidence."
    },
    "1107/1107.3665_arXiv.txt": {
        "abstract": "{The reason for the observed thinness of the solar tachocline is still not well understood. One of the explanations that have been proposed is that a primordial magnetic field renders the rotation uniform in the radiation zone.} {We test here the validity of this magnetic scenario through 3D numerical MHD simulations that encompass both the radiation zone and the convection zone. } {The numerical simulations are performed with the anelastic spherical harmonics (ASH) code. The computational domain extends from $0.07\\;R_\\odot$ to $0.97\\;R_\\odot$.} {In the parameter regime we explored, a dipolar fossil field aligned with the rotation axis cannot remain confined in the radiation zone. When the field lines are allowed to interact with turbulent unstationary convective motions at the base of the convection zone, 3D effects prevent the field confinement.} {In agreement with previous work, we find that a dipolar fossil field, even when it is initially buried deep inside the radiation zone, will spread into the convective zone. According to Ferraro's law of iso-rotation, it then imprints on the radiation zone the latitudinal differential rotation of the convection zone, which is not observed.}  \\date{Received 14 january 2011 / Accepted 12 June 2011} ",
        "introduction": "\\label{sec:introduction} The discovery of the solar tachocline can be considered as one of the great achievements of helioseismology (\\textit{see} \\citealt{Brown:1989p7}). In this layer the rotation switches from differential (i.e., varying with latitude) in the convection zone to nearly uniform in the radiation zone below.  Its extreme thinness (less than 5 \\% of the solar radius) came as a surprise  (\\textit{see} \\citealt{Charbonneau:1999p19}) and has still not  been explained properly.  \\citet{Spiegel:1992p5} made the first attempt to model the tachocline, and they showed that it should spread deep into the radiation zone owing to thermal diffusion. As a result, the differential rotation should extend down to $0.3\\, R_\\odot$ after $4.5$ Gyr, contrary to what is inferred from helioseismology \\citep{Schou:1998p3791,Thompson:2003p2511}. The authors suggested one way to confine the tachocline, namely to counter-balance thermal diffusion by an anisotropic turbulence, which would smooth the differential rotation in latitude. This model was successfully simulated in 2D by \\citet{Elliott:1997p822}.  \\citet{Gough:1998p34} (GM98 hereafter) argued that such an anisotropic turbulent momentum transport is not able to erode a large-scale latitudinal shear, citing examples from geophysical studies (\\textit{see} \\citealt{starr1968physics} and more recently \\citealt{Dritschel:2008p4287} for a comprehensive review). For instance, the quasi-biennal oscillation in the Earth's stratosphere arises from the anti-diffusive behavior of anisotropic turbulence. But the question is hardly settled because \\citet{Miesch:2003p523} found in his numerical simulations that this turbulence is indeed anti-diffusive in the vertical direction, but the shear is reduced in latitude. Similar conclusions were reached through theoretical (analytical) studies by \\citet{Kim:2005p623}, \\citet{Leprovost:2006p549}, and \\citet{Kim:2007p555}.  Gough \\& McIntyre proposed an alternate model, in which a primordial magnetic field is buried in the radiative zone and inhibits the spread of the tachocline. This is achieved through a balance between the confined magnetic field and a meridional flow pervading the base of the convection zone. Such a fossil field was also invoked by \\citet{Rudiger:1997p17} to explain the thinness of the tachocline and by \\citet{Barnes:1999p2302} to interpret the lithium-7 depletion at the solar surface. Moreover, it would easily account for the quasi uniform rotation of the radiative interior, which the hydrodynamic model of Spiegel \\& Zahn does not (at least in its original formulation). However, one may expect that such a field would spread by ohmic diffusion, and that it would eventually connect with the convection zone, thus imposing the differential rotation of that zone on the radiation zone below.  In order to test the magnetic confinement scenario and to put constraints on the existence of such a primordial field, several numerical simulations of the solar radiation zone were performed by \\citet{Garaud:2002p27} in 2D, and by \\citet{Brun:2006p24} (\\textit{i.e.} paper I, hereafter BZ06) in 3D. These simulations exhibit a propagation of the differential rotation into the radiation zone (primarily in the polar region). The primordial magnetic field successfully connects to the base of the convection zone where the shear is imposed and, according to Ferraro's law of iso-rotation (\\citealt{Ferraro:1937p830}), the angular velocity is transmitted along the magnetic field lines from the convection zone into the radiative interior, much faster than it would through thermal diffusion. However, in the aforementioned calculations the base of the convection zone was treated as an impenetrable wall, \\textit{i.e.} radial motions could not penetrate in either direction (from top down or from bottom up). No provision was made for a meridional circulation originating in the convection zone that would penetrate into the radiation zone and could prevent the upward spread of the magnetic field. This penetration was taken into account in numerical simulations by \\citet{Sule:2005p31}, \\citet{Rudiger:2007p40}, \\citet{Garaud:2008p36}, in a model of the polar region by \\citet{Wood:2007p4288}, and in simulations coupling the radiative zone to a convective zone by \\citet{Rogers:2011p1234} to properly represent the GM98 scenario. In most of these models, an axisymmetric and stationary meridional circulation was imposed at the top of the radiation zone, strong enough  to bend the magnetic field lines of the primordial field. As could be expected, various profiles of meridional circulation resulted in different confinement properties, thus emphasizing the need for a more self-consistent approach. We thus propose here to treat the problem by letting a genuine convective envelope dynamically generate its meridional circulation and differential rotation and to study how these nonlinearly generated flows interact with a fossil field.  Our 3D simulations treat the nonlinear coupling of the convection and radiation zones, and the model includes all physical ingredients that play a role in the tachocline confinement (thermal and viscous diffusion, meridional circulation, convective penetration, solar-like stratification, magnetic fields, pumping, ...). The paper is organized as follows. In Sect. \\ref{sec:model} we describe the model we use to address the question of the magnetic confinement of the tachocline. We then emphasize in Sect. \\ref{sec:general-evolution} the global trend of our simulations, and examine in Sect. \\ref{sec:analysis} the dynamics of the tachocline region. Discussions and conclusions are reported in Sect. \\ref{sec:disc-concl}. ",
        "conclusions": "\\label{sec:disc-concl}  The existence of a magnetic field in the solar radiation zone has been the subject of a longstanding debate. Such a field, for example of fossil origin, could easily account for the uniform rotation of the radiative interior. It could also prevent the inward spread of the tachocline, through thermal diffusion, as was proposed by \\citet{Gough:1998p34}. However, it has been argued that such a field would diffuse into the convection zone and imprint the latitude-dependent rotation of that region on the radiation zone below (which is not observed by helioseismology inversions).  This objection has indeed been confirmed by the calculations of \\citet{Garaud:2002p27} and \\citet{Brun:2006p24}. However, their simulations treated the top of the radiation zone as a non-penetrative boundary, which does not allow for a meridional flow that would originate in the convection zone and could oppose the diffusion of the magnetic field.  This restriction was lifted in the $2$D simulations performed by \\citet{Sule:2005p31,Rudiger:2007p40,Garaud:2008p36}, who showed that this meridional circulation could indeed prevent the magnetic field from diffusing into the convection zone. \\mybold{Very recent work in 2D \\citep{Rogers:2011p1234} has studied the role of a fossil magnetic field in a simulation coupling a radiative and a convective zone. It confirms our result that the field is not confined at the base of the convection zone. However, the author does not find any magnetic transport of angular momentum, possibly because of a force-free state for the magnetic field.} These $2$D simulations assumed large-scale axisymmetric flows, unlike the three-dimensional motions of the real convection zone. For this reason we decided to treat the problem using a 3D code, with the best possible representation of the turbulent convective motions.  We showed by \\mybold{considering time-dependent 3D motions} that the interior magnetic field does not stay confined in the radiative interior, but expands into the convection zone. Non-axisymmetric convective motions interact much differently with a dipolar-type magnetic field, as would a monolithic, single cell large-scale meridional circulation.  \\mybold{That our magnetic field is able to expand from the radiative interior into the convection zone results in an efficient outward transport of angular momentum through the tachocline, imposing Ferraro's law of iso-rotation. It is likely a direct consequence of the non force-free configuration of our magnetic field that yields this efficient transport through the large-scale magnetic torque.}  We are aware that our simulations are far from representing the real Sun, because of the enhanced diffusivities we had to use to meet numerical requirements. \\mybold{The overshooting convective motions are overestimated owing to our enhanced thermal diffusivity $\\kappa$. The high solar Peclet number (lower $\\kappa$) regime would make the flow amplitude drop even faster at the base of the convection zone, which would result in an even weaker meridional circulation velocity. This would not help the confinement of the magnetic field. \\\\ Although our enhanced magnetic diffusivity is quite large, we demonstrated (\\textit{see} Sect. \\ref{sec:general-evolution}) that our choice of parameters ensured that we are in the correct regime for magnetic field confinement, with advection dominating diffusion.} Nevertheless, more efforts will be made  in future work to increase the spatial resolution and to improve the subgrid-scales representation.\\\\ \\mybold{The initial magnetic topology may also be questioned. The gradient of our primordial dipolar field possesses a maximum at the equator. Because the field is initialized in the radiation zone where its evolution is diffusive, the magnetic field initially evolves faster at its maximum gradient location, \\textit{i.e.}, at the equator. As a result, it first encounters the complex motions coming from the convection zone at the equator. Other simulations conducted with a modified topology to change the location of the maximum gradient (and thus the latitude where the magnetic field primarily enters the tachocline) led to similar results as those reported here.}  \\mybold{Given the extension of our tachocline}, the differential rotation is already initially present below the base of the convection zone before we \\mybold{introduce} the magnetic field. A confined magnetic field would naturally erode this angular velocity gradient, though we do not know much about the real initial conditions. In our simulation the magnetic field opens into the convection zone, consequently there is no chance to obtain a solid rotator in the radiative zone because there is no confinement. In fact, our magnetic field lines primarily open at the equator. We point out that this is a major \\mybold{difficulty with} the GM98 scenario: even if magnetic pumping can be a way to slow down the penetration of the magnetic field into the convective zone, it will not prevent it from connecting the two zones, and/or prevent the magnetic field lines from opening there.  Furthermore, the meridional circulation is not a laminar \\mybold{uni-cellular} flow in our 3D simulation. It has a complex instantaneous multicellular pattern, which affects the efficiency of a polar confinement.  We therefore conclude that a dipolar axisymmetric fossil magnetic field cannot prevent the spread of the solar tachocline. We intend to explore other mechanisms and different magnetic field topologies. Our 3D model will allow us to test non-axisymmetric configurations, \\textit{e.g.}, an oblique and confined dipole."
    },
    "1107/1107.4103_arXiv.txt": {
        "abstract": "We study the scaling between bulge magnitude and central black hole (BH) mass in galaxies with virial BH masses $\\lesssim 10^6\\msun$. Based on careful image decomposition of a snapshot \\emph{Hubble Space Telescope} $I$-band survey, we found that these BHs are found predominantly in galaxies with pseudobulges.  Here we show that the \\mbulge\\ relation for the pseudobulges at low mass is significantly different from classical bulges with BH masses $\\geq 10^7\\msun$.  Specfically, bulges span a much wider range of bulge luminosity, and on average the luminosity is larger, at fixed $M_{\\rm BH}$.  The trend holds both for the active galaxies from Bentz et al. and the inactive sample of G\\\"ultekin et al. and cannot be explained by differences in stellar populations, as it persists when we use dynamical bulge masses.  Put another way, the ratio between bulge and BH mass is much larger than $\\sim 1000$ for our sample. This is consistent with recent suggestions that $M_{\\rm BH}$ does not scale with the pseudobulge luminosity.  The low-mass scaling relations appear to flatten, consistent with predictions from Volonteri \\& Natarajan for massive seed BHs. ",
        "introduction": "\\label{sec:intro}  For massive elliptical galaxies and galaxies with classical bulges, black hole (BH) masses have been found to correlate with various properties of bulge, such as bulge mass \\citep[e.g.,][]{Magorrianetal1998}, bulge luminosity (the \\mbulge\\ relation, see \\citealt{MarconiHunt2003}), and the velocity dispersion of stars in the bulge (the \\msigma\\ relation, see \\citealt{Tremaineetal2002}).  This has led to theoretical work exploring the importance of AGN feedback on galaxy evolution and the creation of the scaling relations \\citep[e.g.,][]{CiottiOstriker2007}. However, we still do not understand the origin of the scaling relations.  Observationally, they are best-measured using nearby elliptical galaxies with central supermassive BHs having masses $\\gtrsim10^7\\msun$.  It is still unclear what the scaling relations look like for BHs with masses $\\sim 10^5-10^6\\msun$ in late-type galaxies. Some simulations \\citep[e.g.,][]{VolonteriNatarajan2009} show that the \\msigma\\ relation can result from the the hierarchical merging of massive dark matter halos \\citep[also see][]{Peng2007}.  In that picture, and assuming massive BH seeds ($\\sim10^5\\msun$), the \\msigma\\ relation would flatten at the low-mass end. Measurement of the scaling relations of low-mass BHs will directly constrain models of BH seed formation and help elucidate the most important factors in establishing the observed scaling relations.  Based on \\emph{Hubble Space Telescope} (\\emph{HST}) observations of the sample of $18$ low-mass BHs selected from the first data release (DR1) of the Sloan Digital Sky Survey (SDSS; \\citealt{GreeneHo2004}), \\cite{Greenetal2008} study the \\mbulge\\ relation at low mass and find a different scaling, with BHs living in larger bulges than expected from the high-mass relations. However, their sample is small and the result is quite uncertain. \\cite{Greenetal2007} present 174 low-mass BHs from DR4 of the SDSS, giving us the opportunity to study the scaling relations with an order of magnitude more objects. Their \\msigma\\ relation has been examined by \\cite{Xiaoetal2011}. We have done detailed image decompositions of these galaxies based on \\emph{ HST} imaging \\citep[][hereafter Paper I]{Jiangetal2011} and the \\mbulge\\ relation will be studied here.  Low-mass BHs are found in different host galaxy types than more massive BHs.  Most of their host galaxies actually contain pseudobulges.  Pseudobulges have quite different properties from those of violently and rapidly formed classical bulges. For example, pseudobulges have flatter shapes than those of classical bulges (e.g., S{\\'e}rsic index $n<2$).  They have smaller velocity dispersions at fixed luminosity in the Faber-Jackson relation \\citep[][]{FaberJackson1976} compared with classical bulges (e.g., \\citealt{KormendyIllingworth1983,KormendyKennicutt2004}; Paper I). Pseudobulges are believed to form via secular processes in disk galaxies (see the review by \\citealt{KormendyKennicutt2004}). Considering their differing formation mechanisms and properties, we may expect differing BH scaling relations as well, which is observed with small samples with dynamical BH masses (e.g., \\citealt{Hu2008,Hu2009,Greenetal2010}) and indirect analysis (e.g., \\citealt{GadottiKauffmann2009}).  As the bulges hosting the low-mass BHs presented here are likely dominated by pseudobulges (Paper I), studying the properties of these galaxies will also help us understand the growth and evolution of pseudobulges. ",
        "conclusions": "We investigate the \\mbulge\\ and \\mgal\\ relations focused on $10^5 <\\mbh/\\msun< 10^6$.  Both the \\mbulge\\ and \\mgal\\ relations present the same conclusion that these galaxies appear to have smaller BH masses at a given bulge mass than more massive, bulge-dominated systems.  Here we discuss the possible origins of this difference.  It cannot be easily explained by errors in the BH mass scale. We would have to boost the BH masses by $\\sim 1$ dex, which we consider unlikely given that the sample appears to obey an \\msigma\\ relation \\citep[][]{Barthetal2005, GreeneHo2006,Xiaoetal2011}.  Likewise, even correcting for young stellar populations in these bulges, we still find that our BHs are not consistent with the low mass extrapolation of the \\mbulge\\ relation for more massive BHs (consistent with \\citealt{Hu2008,Greenetal2008,GadottiKauffmann2009,Greenetal2010}).  As justified in Paper I, most of our galaxies are consistent with the properties of pseudobulges \\citep[e.g.,][]{KormendyKennicutt2004, FisherDrory2010}.  We believe that pseudobulges are built by secular processes driven by disk instabilites.  If the violent merging that builds classical bulges also establishes BH-bulge scaling relations, we would expect different (or nonexistent) scalings for pseudobulges.  We confirm the conclusion of \\cite{Kormendyetal2011} that there is effectively no correlation between BH mass and pseudobulge luminosity based on our much larger sample. Furthermore, we extend this conclusion to bulge dynamical mass.  This finding suggests that indeed, pseudobulges have quite different properties compared with classical bulges over a wide dynamic range, from $\\sim10^5\\msun$ to $\\sim10^8\\msun$ BHs.  However, our sample is different from the sample of Kormendy et al. in that there is no systematic offset between the Kormendy et al. and the G\\\"ultekin et al. sample.  Thus we see additional tentative evidence for larger systematic differences as a function of mass.  Taking our sample alone, the host galaxies have a wide range of bulge magnitudes over a very limited range in BH mass. In other words, the \\mbulge\\ relation flattens out at the low-mass end.  This seems to be consistent with the results from simulations of \\cite{VolonteriNatarajan2009} for massive BH seed models. If BH seeds can be as massive as $\\sim10^5\\msun$, the low-mass BHs in our sample may be remnants of BH seeds that have not evolved much since they were formed. Consistent with our previous claims, we suggest that secular processes are relatively inefficient at feeding central BHs. Thus bulges that are built secularly maintain a low BH mass independent of the growth of their surrounding bulge."
    },
    "1107/1107.4051_arXiv.txt": {
        "abstract": "We present a multi-facility study of the optical spectrum of the extragalactic globular cluster RZ2109, which hosts a bright black hole X-ray source.  The optical spectrum of RZ2109 shows strong and very broad [O~III]$\\lambda\\lambda$4959,5007 emission in addition to the stellar absorption lines typical of a globular cluster.  We use observations over an extended period of time to constrain the variability of these [O~III] emission lines. We find that the equivalent width of the lines is similar in all of the datasets; the change in L[O~III]$\\lambda$5007 is $\\ltsim 10\\%$ between the first and last observations, which were separated by 467 days.  The velocity profile of the line also shows no significant variability over this interval. Using a simple geometric model we demonstrate that the observed [O~III]$\\lambda$5007 line velocity structure can be described by a two component model with most of the flux contributed by a bipolar conical outflow of about 1,600 \\kms , and the remainder from a Gaussian component with a FWHM of several hundred \\kms. ",
        "introduction": "The first unambiguous black hole X-ray source in a globular cluster was discovered by Maccarone et al.\\ (2007) in the extragalactic globular cluster RZ2109. This globular cluster is in the Virgo elliptical galaxy NGC~4472 (M49), and is a luminous, metal-poor cluster, with an absolute magnitude of M$_V = -10.0$ and a $B$-$V$ color of $0.68$, located $6.6'$ away from the center of its host galaxy \\citep{RZ01}.  The black hole X-ray source in RZ2109 was discovered in XMM observations of the NGC~4472 globular cluster system, of which RZ2109 is a member.  The XMM observations of the source XMMU122939.7+075333 in RZ2109 showed that it had a X-ray luminosity of $\\simeq 4 \\times 10^{39}$ ergs~s$^{-1}$ with x-ray counts that varied by a factor of 7 over a few hours \\citep{Maccarone07}. The high X-ray luminosity, more than an order of magnitude higher than the Eddington luminosity for a neutron star, requires either a black hole or multiple neutron stars in the old population of the globular cluster. The strong variability within the XMM observation rules out multiple neutron stars, and indicates that the source is a black-hole system \\citep{Maccarone07}.  A study of existing optical spectra of RZ2109 revealed that it had broad, strong [O~III]$\\lambda$5007 emission at the absorption line radial velocity of the globular cluster \\citep{Zepf07}.  Follow-up optical spectroscopy with the Keck telescope showed stellar absorption lines typical of an old globular cluster and remarkably broad and luminous [O~III]$\\lambda \\lambda$4959,5007 emission lines \\citep[hereafter Z08]{Zepf08}. Specifically, Z08 showed that the velocity width of the [O~III]$\\lambda$5007 line is $\\simeq 2,000$ km~s$^{-1}$, and that the line luminosity is about $1.4 \\times 10^{37}$ ergs~s$^{-1}$. Moreover, there is no sign of emission lines other than [O~III], and the L([O~III])/L(H$\\beta$) ratio appears to be at least 30 (Z08).   The observed high luminosity and broad velocity width of the [O~III]$\\lambda \\lambda$4959,5007 emission lines place strong constraints on the origin of these emission lines and the nature of the black hole X-ray source in RZ2109. For example, the presence of this very strong and broad [O~III] emission in the black hole X-source hosting globular cluster RZ2109 and the absence of similar emission in any other spectroscopic study of globular clusters argues strongly against significant beaming of the X-rays \\citep{Gnedin2009}.  Furthermore, Z08 showed that the observed broad [O~III]$\\lambda\\lambda$4959,5007 velocities cannot be explained by gravitational motions near the black hole because the available volume sufficiently close to the black hole is many orders of magnitude too small to produce the observed [O~III]$\\lambda$5007, given the critical density of the line. Instead, the broad velocity and high luminosity requires a strong outflow from the accreting black hole.  Z08 suggests that such strong outflows occur for systems near their Eddington limit \\citep[see also][and references therein]{Proga2007}.  The observed L$_{X}$ of $\\simeq 4 \\times 10^{39}$ ergs~s$^{-1}$ thus indicates a stellar mass for the accreting black hole in the globular cluster RZ2109 (see discussion in Z08).  In this study we present several sets of new optical spectroscopic data on the broad [O~III]$\\lambda\\lambda$4959,5007 lines in RZ2109.  We use the multiple observations at different times to study the variability of the [O~III]$\\lambda \\lambda$4959,5007 emission over baselines ranging from one to 467 days. We also take advantage of the higher signal-to-noise and higher spectral resolution of our new Gemini data to constrain models of the [O~III]$\\lambda \\lambda$4959,5007 emitting regions. ",
        "conclusions": "\\label{sec:dc} \\subsection {Variability}  There are two major modes by which the observed lack of strong variability in the [O~III] emission might be produced in an accreting black hole system; a steady accretion mode, or an emitting system in which the crossing time for ionizing photons is large in comparison with the variability timescale of the ionizing source. In the case of a variable ionizing source the emission variability may also contain information on the density distribution of the emitting gas.  The x-ray data described in \\citet{Shih08} span nine years and display a consistent count rate between observations. The four observations described in the present work also display consistent [O~III]$\\lambda$5007 equivalent widths between observations. For reference the light crossing path bounded by the optical observations presented here is 0.4 pc. Both observations are consistent with a long lived steady accretion rate. The sparse temporal sampling of the data, however, does not allow tight constraints on shorter term fluctuations in either the X-ray or [O~III] emission. Future x-ray observations of the RZ2109 x-ray source, XMMU 1229397+07533, with finer temporal resolution could be combined with the [O~III] variability limits presented here to construct limits on the spatial scale of the emitting region.  Likewise variability in the [O~III] emission under constant x-ray flux could yield information on the outflow rate and density distribution.  The absence of strong variability does allow for the elimination of one class of models as the source of the observed [O~III] emission.  Novae shells in the local Milky War are a known source of [O~III]$\\lambda$5007 emission \\citep{Downes2001}. The very brightest shell nova reach L([O~III]$\\lambda$5007) $\\sim 10^{37}$ ergs and have velocity widths on the order of those observed in RZ2109 \\citep[][for example]{Downes2001,Iijima2011}.  In contrast to the observations presented here, however, the [O~III] luminosity of these sources decay by several orders of magnitude and velocity profile vary on timescales of days to weeks.  \\subsection{Structure of Emission Region}  It is somewhat remarkable that the geometric model presented in Section \\ref{sec:gm} fits the [O~III] line profile as well at it does with simplistic assumptions. There are, of course, a number of questions either not addressed or raised by the model, including the spatial scale of the emitting system, the observed flux asymmetry of the red and blue high velocity width component, the failure of the model at the extreme high velocity wings, and the nature of the low velocity Gaussian component.  The geometric model as presented is free of any spatial scaling.  We must look to other means in order to place limits on the size of the outflow region.  A firm lower limit may be set using the total emission from the high velocity width outflow component of the [O~III]$\\lambda 5007$ emission, adopting a pure [O~III] gas at critical density for [O~III]$\\lambda 5007$ at a temperature of $10^{4}$ K \\citep{Osterbrock89}, and employing the geometry parameters found. With this minimal constraint the observed emission could originate from a region with a minimum radius of a few $\\times 10^{-3}$~pc.  The data do not provide a clear upper limit on the size of the emitting region, and there is no reason given the available data that the size of the [O~III]$\\lambda 50007$ region could not be as large as the $\\sim 10$ pc extent of the globular cluster itself.  Asymmetries in the bipolar outflows are a well known feature of accreting black hole systems. When a flux asymmetry is attributed to a differential obscuration the usual pattern observed is a decrement in the flux of the receding flow produced by the additional obscuring material along the greater line of sight distance \\citep{vanG87,Peterson00}. However, we observe the opposite effect in the RZ2109, with the receding [O~III] emission line 40 percent more luminous than the approaching flow.  If differential obscuration contributes to the observed flux asymmetry, the obscuring media must be local to the cluster system, such as dust-bearing gas at large radius resulting from an earlier outflow of the system or interaction with a asymmetric ISM. The mechanism by which such a local obscuring media may be produced is unclear.  A second known source of flux asymmetry in outflows is Doppler boosting \\citep{MR99,Fender2006}. In Doppler boosting, however, the flux of the approaching outflow is amplified relative to the receding outflow, again opposite of the asymmetry observed in RZ2109.  The different line of sight travel times of outflows toward and away from the observer can also produce asymmetric fluxes if coupled with a variable ionizing source.  However, the flux asymmetry is the same in both our Keck and Gemini observations separated by more than a year, which would be very difficult to understand in such a scenario.  It therefore seems most likely that the somewhat greater flux on the red side of the [O~III]$\\lambda 4959$,$\\lambda 5007$ line profile relative to the blue side is a result of an intrinsic asymmetry in the source's emission density. Such an asymmetry may be attributable to either a difference in the mass of each outflow or the ionizing flux incident on them.   An expanding gas distribution such as nova shell or other remnant produced by late stage stellar evolution and ionized by an external source may produce a flux asymmetry across the velocity profile provided the gas is optically thick. Such an arrangement would allow for a flux excess of the receding side of the velocity profile relative to the approaching if the media was positioned between the observer and the ionizing source.  For example, a late stage nova shell modeled with a uniform density shell illuminated by an external x-ray source produces an asymmetric velocity profile with a continuous blue to red flux gradient whose slope is set by the position of the external x-ray source. However, this type of flux asymmetry fails to provide an adequate description of flat topped step function profile observed in the [O~III]$\\lambda 4959$,$\\lambda 5007$ line profiles high velocity width component.  The failure of the model at the most extreme velocities, seen in velocities in excess of $\\pm$1400 km~s$^{-1}$ in Figure \\ref{fig:mod} is most likely attributable to the simplifying assumption of constant bulk velocity in the geometric model. If this is the case and the emission in these extreme velocity wings is produced by the outflow closest to the source, prior to any deceleration, with more sophisticated modeling techniques the extreme velocity wings may provide another constraint on the spatial scale of the system.  While a model with a strong bipolar outflow can naturally account for broad [O~III]$\\lambda 4959$,$\\lambda 5007$ emission, the origin of the narrower Gaussian component is less clear. The observed FWHM of this lower velocity width Gaussian peak is three times the FWHM of the spectral resolution of the Gemini spectrum. Thus, the observed strongly peaked [O~III]$\\lambda 5007$ line profile is representative of its intrinsic shape. One implication of the strongly peaked line profile is that rotation is ruled out for the origin of the observed width of the line. It is worth noting that forbidden lines such as [O~III] are optically thin, and the observed profile reflects the intrinsic one, unlike the much more complex case of optically thick permitted lines such as Balmer emission lines in AGN.  Therefore, for optically thin [O~III], rotation naturally produces a double-peaked or flat-topped velocity profile \\citep[e.g.][]{Clark79}, and fails badly to match the centrally peaked profile seen in the [O~III] emission. For example, the specific case of a rotating Keplerian disk with an inverse square law emission density profile has a velocity profile which deviates from the observed lower velocity width [O~III] component at the $7\\sigma$ level in a one spectral resolution element bin about line center.  In fact, instead of showing a double-horned profile or very flat-topped profile typical of rotation, the observed profile is more centrally peaked than a Gaussian at about the $3\\sigma$ level.  Therefore, rotation is ruled out as the source of the velocity width of the lower velocity width Gaussian component, contrary to the assumption of some modeling \\citep{Porter2010}.  The luminosity of the L([O~III])$\\lambda 5007$ at a given velocity can also be used to determine whether the velocity width of the line can be due to gravitational motions or not.  Specifically, because [O~III]$\\lambda 5007$ is a forbidden line with a known critical density \\citep{Osterbrock89}, we can determine the maximum possible L([O~III])$\\lambda 5007$ within a given volume by adopting a pure [O~III] gas of uniform density at its critical density. We can also compute the largest volume available sufficiently close to the black hole to account for the observed velocity width due to gravitational motions about the black hole.  \\citet{Zepf08} carried out this calculation for the broad velocity component of the [O~III]$\\lambda 5007$, and showed it is many orders of magnitude stronger than can be accounted for with gravitational motions, and to thus require strongly outflowing material.  Here we apply this same calculation to the lower velocity width Gaussian component which is seen much more clearly in our new Gemini data with its greater spectral resolution and higher signal-to-noise. Adopting a pure [O~III] gas of uniform density at its critical density \\citep{Osterbrock89} and a canonical black hole mass of 10\\Msun, we find the maximum luminosity in a one resolution element bin -170 km s$^{-1}$ from line center is of order 10$^{29}$ ergs s$^{-1}$. In contrast, the observed luminosity at this velocity is $4.2\\times 10^{35}$ ergs s$^{-1}$, and thus gravitational motions also fail to account for the velocity width of the lower velocity width Gaussian peak.  It is worth noting that the above arguments do not strictly rule out rotation of the emission line region, they do, however, require that the observed line widths be produced by some mechanism other than rotation. A scenario where the emitting region rotates in the plane of the sky, or with specific arrangements of absorbing systems may be invoked, but in these cases the line width and profile would be independent of any rotation.  Because the emitting region is taken to be optically thin to the forbidden line emission observed, self absorption or intermediate line absorption systems are unlikely. Gray body or broad band absorption/scattering via an obscuring media is possible, however the specific distribution of the material necessary to transform a rotational line profile to that observed (a preferential and symmetric obscuration of the emitting gas moving parallel/anti-parallel with the line of sight) is difficult at best.  Although the specific calculation above is for a rotating sphere for simplicity, adopting a dynamically hot model for the emitting material will change the gravitational motions prediction far less than the five or more orders of magnitude difference between it and the data.  The calculation also adopts a stellar mass for the black hole in the globular cluster, based on results in Z08. In that work, Z08 used the same type of volume calculation as above for the broad component of the [O~III]$\\lambda 5007$ emission, and showed that it is not due to gravitational motions unless the black hole is nearly as massive as the $\\sim 10^6 \\Msun$ cluster itself \\citep[for a more detailed treatment see also][]{Porter2010}. The broad component then must be due to a strong outflow.  Z08 suggest that observational evidence and theoretical calculations \\citep[and references therein]{Proga2007} indicate that strong outflows are associated with high (L/L$_{\\rm Edd}$) systems, which given the observed L$_X$, implies a stellar mass for the black hole in RZ2109.  An alternative idea, that the source could be a tidal disruption or detonation of a white dwarf by an intermediate-mass black hole, fails on several grounds \\citep[e.g.][]{Irwin2010}.  A disruption is unable to account for either the high L([O~III]$\\lambda 5007$) over an extended period of time or the long duration of its high X-ray flux from at least the early 1990s ROSAT era through 2004 \\citep{Maccarone07,Shih08}.  A detonation may account for the high velocity outflow, but does not provide an explanation for why only oxygen lines are seen in the spectrum when the detonation is powered by the fusion of oxygen into iron. Furthermore, neither a disruption or detonation is able to account for the more than a decade of consistently high L$_X$ followed by the recently observed steep decline of L$_X$ by more than an order of magnitude \\citep{Maccarone10}.  With the current observations we have limited ability to address the energetics of the outflow. Since we have no direct measurement of the temperature, we can estimate the mass of [O~III] in the emitting region finding using a temperature typical of [O~III] emission line. With an assumption of $\\rm{T}=10^4$ K we find an [O~III] mass of $7.3 \\times 10^{-5}$ $\\rm{M_{\\sun}}$.  For temperatures from $5 \\times 10^3$--$10^6$ K the mass ranges from $1 \\times 10^{-3}$ to $4 \\times 10^{-6}$ $\\rm{M_{\\sun}}$.  Unfortunately since we have little to constrain either the composition of gas, beyond that it is likely hydrogen depleted, or the ionization fraction as no [O~II] lines are observed above the detection threshold, we are unable to provide a concrete estimate of the total gas mass of the emitting region. For a pure oxygen gas the total mass will be within a factor of a few greater than the [O~III] mass, while a solar composition would imply a total mass roughly two orders of magnitude greater. Using the 1600 \\kms velocity of the higher velocity component found in Section \\ref{sec:gm} the kinetic energy of the [O~III] is then 10$^{45}$ erg. This energy is 8 orders of magnitude smaller than the available total accretion energy, assuming an accretion efficiency of 10 percent and a 1 M$_{\\sun}$ donor. If a solar composition is assumed for the outflowing gas and all oxygen being doubly ionized the systems kinetic energy is 6 orders of magnitude smaller than the total accretion energy. It should be noted that the kinetic energy given here is for the total observed outflow; it is unknown what fraction of the total accretion event which the current outflow represents."
    },
    "1107/1107.4098_arXiv.txt": {
        "abstract": "Dark matter halo merger trees are now routinely extracted from cosmological simulations of structure formation. These trees are frequently used as inputs to semi-analytic models of galaxy formation to provide the backbone within which galaxy formation takes place. By necessity, these merger trees are constructed from a finite set of discrete ``snapshots'' of the N-body simulation and so have a limited temporal resolution. To date, there has been little consideration of how this temporal resolution affects the properties of galaxies formed within these trees. In particular, the question of how many snapshots are needed to achieve convergence in galaxy properties has not be answered. Therefore, we study the convergence in the stellar and total baryonic masses of galaxies, distribution of merger times, stellar mass functions and star formation rates in the \\glc\\ model of galaxy formation as a function of the number of ``snapshot'' times used to represent dark matter halo merger trees. When utilizing snapshots between $z=20$ and $z=0$, we find that at least 128 snapshots are required to achieve convergence to within 5\\% for galaxy masses. This convergence is obtained for mean quantities averaged over large samples of galaxies---significant variance for individual galaxies remains even when using very large numbers of snapshots. We find only weak dependence of the rate of convergence on the distribution of snapshots in time---snapshots spaced uniformly in the expansion factor, uniformly in the logarithm of expansion factor or uniformly in the logarithm of critical overdensity for collapse work equally well in almost all cases. We provide input parameters to \\glc\\ which allow this type of convergence study to be tuned to other simulations and to be carried out for other galaxy properties. ",
        "introduction": "Simulations of the cosmological evolution of large scale structure and individual galaxies in cold dark matter cosmogonies \\citep{springel_simulations_2005,kuhlen_via_2008,springel_aquarius_2008,boylan-kolchin_resolving_2009,klypin_halos_2010,prada_halo_2011} are an invaluable tool in studying the formation and evolution of galaxies. With the high resolutions and large volumes attained by modern simulations, a particularly useful approach is to extract information on individual dark matter halos and how those halos merge during the process of hierarchical structure formation. It is within these merging hierarchies of dark matter halos that galaxies form. Typically, this information on hierarchical growth is encoded as a ``merger tree'' represented by a set of ``nodes'' (each corresponding to a dark matter halo at one point in its evolution) and ``branches'' connecting those nodes which link descendant halos to their progenitors (i.e. halos which will merge together to form the descendant). While these merger trees have many applications, a very common use is an input to models which aim to solve the physics of galaxy formation within the merging hierarchy of dark matter---in particular the class known as ``semi-analytic models'' \\citep{cole_hierarchical_2000,hatton_galics-_2003,croton_many_2006,de_lucia_hierarchical_2007,monaco_morgana_2007,somerville_semi-analytic_2008}.  Due to the way N-body simulations are analyzed and their data stored, merger trees are usually constructed from a set of ``snapshots''---outputs of all of the particle data at a set of $N$ times, $t_i$ where $i=1$ to $N$. Since $N$ is finite (and often limited by available storage and processing power) this means that each merger tree is a temporally sparse representation of some underlying true merger tree which may have structure on shorter timescales. Such structure is lost in the sparse representation. The question of how this loss of information affects the properties of galaxies produced by models utilizing sparse merger trees has been touched upon before, but never carefully studied. For example, \\cite{helly_galaxy_2003}, utilizing merger trees with 44 snapshots (spaced uniformly in the logarithm of expansion factor) from $z=0$ to $z=20$, compared their results to those obtained from their standard {\\sc Galform} model based on merger trees constructed from Press-Schechter techniques using a significantly larger number of snapshots and noted that ``\\ldots we find that if we degrade the time resolution of the standard {\\sc Galform} model to match that of our N-body model the properties of the galaxy populations predicted change very little.'' However, this statement refers only to statistical properties of the entire galaxy population rather than to individual galaxies and, in any case, no details were given. \\cite{hatton_galics-_2003} utilized merger trees with around 70 snapshots between $z=0$ and $z=10$. They repeated their calculations using only every second snapshot and found that this leads to a 20\\% scatter in cold gas content in galaxies, and a 40\\% scatter in the mass of hot gas associated with each galaxy. They concluded that this was not a significant problem, but this is arguably a significant discrepancy given the accuracy of modern observational datasets. \\cite{croton_many_2006} describe a galaxy formation model utilizing merger trees extracted from the Millennium Simulation with 60 snapshots between $z=0$ and $z=20$ but do not discuss if this number is sufficient to achieve converged results. \\cite{somerville_semi-analytic_2008} present a model based on merger trees constructed using an extended Press-Schechter algorithm and note that they find no significant changes if they instead implement their model in merger trees extracted from N-body simulations, but do not provide further details.  Importantly, none of these studies were able to address the issue of how fast galaxy properties converge as the number of snapshots is increased, or what is the optimal distribution of snapshots in time. Additionally, in all of the galaxy formation models employed at least some of the baryonic physics (e.g. the rate of gas cooling) is solved using timesteps tied to the snapshot spacing---this is non-optimal as it is in principle possible to obtain better convergence by solving the baryonic physics on a much shorter timescale (and interpolating the dark matter halo properties as necessary).  In this work, we address these issues by utilizing the \\glc\\ semi-analytic galaxy formation code \\citep{benson_galacticus:_2010}. This code has two key features which make it ideally suited to this problem. The first is its ability to construct merger trees using a modified extended Press-Schechter formalism \\citep{parkinson_generating_2008} which both agree with the statistics of merger trees extracted from N-body simulations and which can be built without reference to any fixed grid of snapshot times, thereby achieving arbitrarily high temporal resolution. As we will describe below, \\glc\\ is able to post-process these trees to construct sparse representations in the same manner as trees extracted from an N-body simulation. The second key feature of \\glc\\ is that all baryonic physics is solved using adaptive timesteps which are adjusted to keep a specified tolerance in the quantities being computed. As such, the number and distribution of snapshots in the merger trees affects the results from \\glc\\ only due to the inevitable information loss resulting from the sparse representation. We exploit these features to explore convergence of basic galaxy properties as a function of the number of snapshots available and the distribution of those snapshots in time.  The remainder of this paper is arranged as follows. In \\S\\ref{sec:Method} we describe the process by which we construct merger trees and alter their temporal resolution. In \\S\\ref{sec:Results} we present results from our convergence study and in \\S\\ref{sec:Conclusions} we give our conclusions. ",
        "conclusions": "\\label{sec:Conclusions}  It has become common practice to extract histories of the hierarchical merging process (``merger trees'') from cosmological N-body simulations and to use these as inputs to semi-analytic models of galaxy formation. Previously, there has been little consideration of how the temporal resolution of these trees affect the properties of the resulting galaxies. In this work we performed a convergence study using the \\glc\\ toolkit.  We find that 128 snapshots spaced uniformly in the logarithm of expansion factor (or, almost equivalently, in the logarithm of the critical overdensity for collapse) provide good ($\\sim 10\\%$) convergence in galaxy stellar and total masses, the number of viable subhalos (i.e. those which have progenitors in the merger tree that are isolated, non-substructure halos), distributions of merger times, in stellar mass functions at $z=0$ and in the volume density of star formation rate as a function of redshift. Smaller numbers of snapshots lead to rapidly diverging results and should be avoided. We also considered snapshots spaced uniformly in expansion factor. We find that no substantial difference in the number of timesteps required to reach a given degree of convergence using this distribution of snapshots. This convergence is obtained for mean quantities averaged over large samples of galaxies---the full model in particular shows significant variance for individual galaxies even when using very large numbers of snapshots.  Our results should provide guidance as to how many snapshots should ideally be stored from future N-body simulations to ensure that the resulting temporally sparse merger trees do not overly limit the accuracy of subsequent galaxy formation calculations. Our results are for a specific set of cosmological parameters and tree mass resolution, in addition to being for a specific implementation of baryonic physics. The rate of convergence plausibly depends on all of these factors, and will likely differ for galaxy properties other than those considered here. Since \\glc\\ is freely available as an open source project\\footnote{\\protect\\glc\\ can be downloaded from \\href{http://sites.google.com/site/galacticusmodel}{{\\tt http://sites.google.com/site/galacticusmodel}}.}, it is relatively easy for anyone to repeat the analysis performed here for a specific set of simulation parameters and galaxy properties. Our results were obtained with v0.9.0.r491 of \\glc\\ and we have made the input parameters available online\\footnote{The input parameter files and scripts to construct plots of the results can be downloaded from \\href{http://www.ctcp.caltech.edu/galacticus/parameters/dmTreeConvergence.tar.bz2}{{\\tt http://www.ctcp.caltech.edu/galacticus/parameters\\\\/dmTreeConvergence.tar.bz2}}.}.  Increasing mass resolution in simulations implies that merger trees contain more information. However, for this information to be folded into semi-analytic model predictions, trees must be built with finer time-stepping. Thus, increasing mass resolution would imply increasing both the size of each snapshot and the number of such snapshots, thereby causing a ``data tsunami''. It would then be recommendable for large high resolution simulations to be post-processed on-the-fly, writing at finely spaced times only (sub)halo catalogues instead of the entire snapshot.  Recent interest in exploring and constraining the parameter space of semi-analytic galaxy formation models \\citep{henriques_monte_2009,bower_parameter_2010,lu_bayesian_2010} makes it crucial to understand and control numerical inaccuracies in such codes. Otherwise, quantitative constraints on model parameters will be subject to unknown systematic biases. While many uncertainties remain in our understanding of the physics of galaxy formation, it is important to ensure that numerical results are converged. Considerations such as those described here should become a standard part of any galaxy formation study."
    },
    "1107/1107.2441_arXiv.txt": {
        "abstract": "Super-hot looptop sources, detected in some large solar flares, are compact sources of HXR emission with spectra matching thermal electron populations exceeding 30 megakelvins.  High observed emission measure, as well as inference of electron thermalization within the small source region, both provide evidence of high densities at the looptop; typically more than an order of magnitude above ambient.   Where some investigators have suggested such density enhancement results from a rapid enhancement in the magnetic field strength, we propose an alternative model, based on Petschek reconnection, whereby looptop plasma is heated and compressed by slow magnetosonic shocks generated self-consistently through flux retraction following reconnection.  Under steady conditions such shocks can enhance density by no more than a factor of four.  These steady shock relations (Rankine-Hugoniot relations) turn out to be inapplicable to Petschek's model owing to transient effects of thermal conduction.  The actual density enhancement can in fact exceed a factor of ten over the entire reconnection outflow.  An ensemble of flux tubes retracting following reconnection at an ensemble of distinct sites will have a collective emission measure proportional to the rate of flux tube production.  This rate, distinct from the local reconnection rate within a single tube, can be measured separately through flare ribbon motion.  Typical flux transfer rates and loop parameters yield emission measures comparable to those observed in super-hot sources. ",
        "introduction": "Solar flares are best known for releasing large amounts of magnetic energy.  They are equally remarkable, however, for large increases in plasma density accompanying the energy release.  During a large flare the coronal density can increase by two or three orders of magnitude, from an ambient value like $n_e\\sim10^9\\,{\\rm cm}^{-3}$, to $n_e\\sim10^{11}$ or $10^{12}\\,{\\rm cm}^{-3}$ \\citep[see eg.][]{Moore1980b}.  The temperature of this dense plasma is typically one order of magnitude larger than ambient conditions ($T_{\\rm fl}\\sim20$ MK) so the magnetic field must confine plasma pressure three to four orders of magnitude above normal.  The coronal magnetic field is strong enough to confine this high-pressure plasma in the two directions perpendicular to field lines, but the plasma is unconfined along the field lines.  It is commonly believed that transport of mass and energy along the field lines drives a process known as chromospheric evaporation to supply the extra mass responsible for the density enhancement \\citep{Antiochos1978}.  Energy is added to the chromosphere causing it to expand into the corona, thereby diluting the coronal energy over more mass, lowering its temperature but raising its density.  In addition to hot loops, filled by chromospheric evaporation, some flares include an over-dense ``knot'' of plasma at the loop apex known as {\\em looptop source} \\citep[LTS, see review by][]{Fletcher1999}.  It is possible that many more flares have LTSs that cannot be imaged by current methods limited in dynamic range \\citep{Petrosian2002}. An LTS is generally assigned a high plasma density due to its radiative output in hard X-rays (HXRs) by which the sources are most commonly imaged.  Its location atop the loop\\footnote{Some investigators prefer the term ``above-the-looptop'' in recognition of its separation from loops visible in soft X-ray or EUV images.  In all likelihood the LTS is {\\em at the top of} some magnetic loop, even if that loop is not (yet) visible at longer wavelengths.}  suggests the LTS is a direct manifestation of magnetic energy release \\citep{Masuda1994}.  If they are thus assumed to arise early in the energy release sequence, it would seem problematic to invoke evaporation as the source of excess density in LTSs.  Moreover, the source appears physically separated from the chromosphere \\citep{Jiang2006}, suggesting its mass is not supplied from there.  A subset of LTSs appear to be dense enough that their HXR emission arises from an electron population which is convincingly {\\em thermal} (i.e.\\ Maxwellian).  The temperatures of these {\\em super-hot} (SH) sources typically exceed $30$ MK and sometimes reach $50$ MK \\citep{Lin1981,Lin1985,Kosugi1994,Nitta1997}.  One indication of their density comes from the emission measure (EM) of the thermal brehmsstrahlung, typically exceeding $10^{48}\\,{\\rm cm}^{-3}$.  Attributing a large filling factor (often unity) to the apparent source provides density values as high as $n_e\\sim 10^{11}$ \\citep[a smaller filling factor would demand an even higher density]{Veronig2006,Longcope2010,Caspi2010}.  Independent evidence for these high densities comes from the inference of a Maxwellian electron distribution, supported at high confidence by high resolution HXR spectra \\citep{Lin1981,Longcope2010,Caspi2010}. Electrons must collide with one another numerous times to fully thermalize.  Accommodating this within a source extending only $10$ Mm requires electrons with mean free paths $\\ell_{e}\\la1$ Mm along the field line.  Electrons at a typical SH temperature of $T_e=35$ MK  have a collisional mean free path, \\be \\ell_{e} ~\\simeq~ 100\\,{\\rm Mm}\\left({n_e\\over 10^9\\,{\\rm cm}^{-3}}\\right)^{-1} ~~, \\ee along the magnetic field line, so thermalization within the LTS demands $n_e\\ga10^{11}\\,{\\rm cm}^{-3}$.  Generally, LTSs show a {\\em non-thermal} electron population ({\\em e.g.}\\ a power-law), similar to footpoint sources, during a flare's impulsive phase \\citep{Alexander1997}.  Superhot thermal sources typically appear {\\em after} the impulsive phase and persist $\\sim10$ minutes or more, far longer than the conductive cooling time \\citep{Jiang2006}.   Some HXR spectral analyses show the SH source co-existing with a second thermal population at $T_e\\sim 20$ MK and much higher emission measure \\citep[5--20 times higher,][]{Longcope2010,Caspi2010}.  This lower-temperature component closely matches the soft X-ray light-curve, from GOES for example, commonly attributed to chromospheric evaporation.  This re-affirms a role for the SH component as a flare stage {\\em prior} to evaporation, thus demanding a different mechanism for enhancing its density.  To date, little effort has been devoted theoretically explaining the high densities in LTSs.  Non-thermal sources are sometimes assumed to be filled by chromospheric evaporation driven by impact of precipitating particles \\citep{Veronig2004}.  Doing so tacitly assumes an energy release process able to persist long enough {\\em on a single field line} to interact with its own effects: chromospheric evaporation.  This seems at odds with the favored energy release mechanism, magnetic reconnection, whereby a single field line will reconnect once in  a single instant.  Resolution of this puzzle seems to require a self-consistently coupled model of particle acceleration and magnetic reconnection; such a model is still still being pursued.  Since SH looptop sources are thermal, and show little evidence for non-thermal particles, they might be understood without an understanding of particle acceleration.  Most widely accepted models of magnetic reconnection, such as that of \\citet{Petschek1964}, use fluid equations and are therefore directly applicable to SH sources.  Several LTS models proposed so far have invoked an increase in magnetic fields strength, a ``collapsing trap'', to enhance the density and temperature \\citep{Somov1997,Karlicky2004,Caspi2010}.   To match observed densities these often require field strengths to increase by factors of 30--100.   These models were not based on reconnection solutions, in fact, assuming a field strength {\\em increase} would seem to contradict spontaneous reconnection scenarios where energy is released by {\\em reducing} magnetic energy.  The notable exception is the fast magnetosonic termination shock which sometimes forms where the reconnection outflow jet encounters an ``obstacle'', such as the arcade of previously reconnected loops.  Fast magnetosonic termination shocks have been a common feature in reconnection-driven flare models \\citep{Forbes1983,Forbes1986,Tsuneta1997,Aurass2002,Vrsnak2005}.  The magnetic field strength increases across a shock of this type, making it suitable for Fermi particle acceleration \\citep{Tsuneta1998}.  The density also increases across the shock, but typically by no more than a factor of two \\citep{Forbes1986,Vrsnak2005}.  This modest density enhancement (far short of the factor of 100 often seen in loop-top sources) coupled with the low emission measure expected in a small structure, makes the fast magnetosonic termination shock an unlikely candidate for a collapsing trap or a LTS.  The principal element in Petschek's fast reconnection model is a {\\em slow magnetosonic shock}, or slow shock, which heats and accelerates the plasma.\\footnote{Petschek's 1964 paper contains several pioneering elements which were  extensively studied in subsequent investigations under the generic term ``Petschek reconnection''. It derived a steady external solution matched to a resistive internal solution to obtain its well-known reconnection rate.   This solution was improved and generalized over following decades \\citep{Sonnerup1970,Vasyliunas1975,Soward1982,Priest1986}.  The solution is, however, modified when the electric field is not resistive or is unsteady \\citep{Heyn1996,Nitta2001}.  Petschek's work was also the first to recognize the generation of slow shocks by reconnection, which turn out to be inevitable in most fast reconnection modes \\citep{Semenov1983,Erkaev2000}.  It is to this aspect of Petschek's work which we exclusively refer hereafter.} Since field strength {\\em decreases} across shocks of this type they are unsuitable for Fermi acceleration.  They do, however, compress the plasma considerably more than the fast magnetosoic termination shock can.  Reconnection between {\\em anti-parallel} fields generates slow shocks at the switch-off limit, wherein the maximum possible density enhancement is 2.5 \\citep{Forbes1989}.  Reconnecting skewed fields (i.e.\\ any angle other than anti-parallel) can, however, result in density enhancements as great as four, according to standard shock models.  This compressed and heated plasma forms a long jet extending away from the reconnection site which may have considerably greater emission measure than the small fast magnetosoic termination shock.  For these reasons it seems plausible for LTSs to be manifestations of the plasma heated and compressed by slow shocks according to Petschek's reconnection model.  Recently \\citet{Longcope2010} modeled a SH-LTS observed by RHESSI using slow shocks from a 3-dimensional time-dependent reconnection theory.  The post-shock temperature follows directly from the angle between reconnecting fields lines, which follows in turn from a model of the pre-flare magnetic field.  \\citet{Longcope2010} found the observed SH temperature $T\\simeq40$ MK to be consistent with the angle from their pre-flare magnetic model: $\\Delta\\theta=70^{\\circ}$ (anti-parallel fields have $\\Delta\\theta=180^{\\circ}$).  The time-dependent model assumed reconnection occurred sporadically in numerous small {\\em patches} within the current sheet separating the skewed fields.  Each patch produced a single flux tube whose subsequent retraction through the sheet generated a plug of SH plasma between traveling slow shocks.  Each plug emitted for $\\sim 8$ sec before confining inflows ceased and its own pressure ``disassembled'' it.  The single observed LTS was actually composed of $\\sim30$ separate, unresolved plugs in flux tubes created by separate patches across the sheet.  The plug collection had a total emission measure proportional to the rate of patch-production which was in turn proportional to the mean reconnection rate.  The latter was measured using the motion of chromospheric flare ribbons and found to be consistent with the emission measure of the observed LTS.  Each reconnected flux tube was later observed at lower temperature as a distinct post-flare loop in TRACE 171\\AA\\ images.  The rate of loop appearances was also deemed consistent with the the inferred patch-production rate.  While the compression ratio is greater for a slow shock than for the fast magnetosoic termination shock (4 {\\em vs.} 2), it is still significantly below the level inferred for observations of looptop sources: 10--100.  This fact forced \\citet{Longcope2010} to assume a very high ambient density ($n_{e0}\\simeq 8\\times10^{10}\\,{\\rm cm}^{-3}$) in order to match the observed emission measure.  This same problem faces all models invoking shocks to compress plasma.  The aforementioned  compression ratios, used by all analytical reconnection models, come from standard MHD shock models, based on conservation laws across a steady-state jump \\citep[see][for example]{Priest2000}.  One possible exception occurs when radiation can cool the plasma faster than compression heats it --- so called {\\em radiative shocks} \\citep{Xu1992}.  Unfortunately, the radiative loss function is so low at SH temperatures that even a density of $n_e=10^{12}\\,{\\rm cm}^{-3}$ cannot produce effective cooling ($\\tau_{\\rm rad}\\simeq8$ minutes!).  Thus it would seem impossible to invoke a radiative shock in a SH looptop source.  Shocks generated by magnetic reconnection are, however, so far from steady state conditions that standard compression ratios, so often used, are not actually applicable.  Thermal conduction along the magnetic field line will generate a {\\em heat front} ahead of the shock \\citep{Thomas1944,Grad1951,Germain1960,Forbes1986b,Kennel1988,Yokoyama1997}.  Since thermal conduction does not affect net energy conservation, the heat front does not change the steady-state density ratio. Under coronal conditions thermal conductivity is extremely large and the heat front would be enormous by the time it achieved steady state --  far larger than any conceivable flare loop \\citep{Guidoni2010}.  This forces the conclusion that solar flare shocks never achieve their steady state.  Recent simulations of these shocks by \\citet{Guidoni2010} show density enhancements by factors 6--8 during the transient phase following reconnection.  More careful study, herein presented, reveals there is no upper bound to the transient density enhancements and that factors of 10--100 are reasonable for a large solar flare.  We therefore propose that the large plasma densities in LTSs are the direct result of compression in slow shocks of the fast reconnection liberating the energy.  One advantage of the model here proposed is that it directly couples the density enhancement to the energy release.  In models of fast magnetic reconnection, magnetic energy is liberated primarily by shortening field lines following topological change within a very small diffusion region.  Except in cases of purely anti-parallel reconnection, the field strength decreases only slightly and thus the release cannot be interpreted as {\\em field annihilation}.  The field line shortens at the local Alfv\\'en speed which is far larger than the hydrodynamic sound speed ($\\beta\\ll1$).  The plasma within the shortening flux tube is therefore compressed at {\\em super-sonic} speeds, generating the slow shocks which heat the plasma.  This is a time-dependent process on any given field line, so the shock behavior is transient rather than steady.  Compression and heating, i.e.\\ the slow shock, is thus a primary effect of releasing magnetic energy, while the fast magnetosoic termination shock is a secondary effect caused by abruptly {\\em halting} the energy release.  We propose here that LTSs are manifestations of the primary effect: slow shocks.  We present our model by first reviewing (in \\S 2) how magnetic energy is converted to heat when fast reconnection occurs at a current sheet.  We show that similar predictions are made by two-dimensional steady-state models, such as Petschek's, as by transient patchy reconnection models.  The latter can be modeled as a simple shock tube problem.  In \\S 3 we use the shock-tube problem to demonstrate how thermal conduction leads to density enhancement far beyond what Rankine-Hugoniot relations predict.  We find a simple relation between the shock-tube mach number and the maximum density enhancement.  We next use thin flux tube models to investigate how much density enhancement could occur within a reconnecting current sheet.  In \\S 5 we apply these single-flux tube results to a flare consisting of numerous flux transfer episodes.  We show how the $EM$ of the SH-LTS scales with the rate of flux transfer.   The constant of proportionality is determined by the angle between field lines reconnected at the current sheet.  Section 6 then considers how the energy released by reconnection drives chromospheric evaporation to generate a second, cooler component of the plasma.  Thus  while the chromosphere contributes most the the $EM$ in a flare, it is not responsible for the LTS. ",
        "conclusions": "We have herein attempted to show that compression by slow mode shocks, first predicted in the \\citet{Petschek1964} model of fast reconnection, is capable of producing high-density, super-hot loop-top emission like that seen in flares.  The density enhancement achieved in these shocks is considerably higher than generally assumed on the basis of steady shock models (Rankine-Hugoniot relations).  This is due to the cooling of the post-shock plasma by thermal conduction during the initial (transient) phase of the post-reconnection retraction.  This effect has been explored using solutions of one-dimensional shock tube equations and two-dimensional thin flux tube equations.  The simplified shock-tube model led us to an estimate of density enhancement, \\eq\\ (\\ref{eq:rho_ad}), in terms of properties of the current sheet and its plasma.  Figure \\ref{fig:bsweep} uses this relation to show the density and temperature enhancements {\\em vs.} the strength and angle of the fields separated by the current sheet.  Density enhancements as large as 100 appear plausible for the strongest fields ($B\\ga500$ G) and largest angles ($\\Delta\\theta\\ga120^{\\circ}$).  This conforms to observational evidence that SH-LTSs occur only when field strengths are extremely high \\citep{Caspi2010}.  \\begin{figure}[htb] \\epsscale{0.8} \\plotone{f12.ps} \\caption{Estimates of density enhancement over a range of flare parameters.  For ambient density $n_{e0}=3\\times10^{9}\\,{\\rm cm}^{-3}$ and temperature $T_0=2$ Mk, a range of fields strengths (abscissa) and reconnection angles are used evaluate \\eq\\ (\\ref{eq:rho_ad}), (bottom curve).  The temperature at time of this peak value is plotted in the top panel.  The result from simulation in \\fig\\ \\ref{fig:deft_sweep} ($z_R=h/2$) is shown with a triangle.  The lower triangle is the observed enhancement and the upper triangle (on the curve) is the adiabatic value.} \\label{fig:bsweep} \\end{figure}  The shock tube model is a natural abstraction of the thin flux tube model for dynamical evolution following patchy reconnection.  These simplified, time-dependent dynamics turn out to capture the behavior of even classical steady-state MHD reconnection models, provided the reconnecting fields are not perfectly anti-parallel.  In all cases shocks, and corresponding density enhancements, are direct consequences of energy release.  Magnetic energy is released by shortening field lines, which naturally compresses the plasma they link.  Since the field lines shorten at Alfv\\'enic speed the plasma is compressed supersonically ($\\beta\\ll1$) leading to very strong shocks.  There is quantitative agreement on shock-produced density enhancement between analytic treatments which are steady-state \\citep{Soward1982b,Vrsnak2005} and transient \\citep{Lin1994,Longcope2009}.  This agreement occurs because even in steady-state models, the response of any single field line to an instantaneous topology change is necessarily transient.  All such models have been based on Rankine-Hugoniot jump conditions across the shocks.  We have here shown these steady-shock Rankine-Hugoniot relations are inapplicable to the transient model since a heat front of asymptotic form is never achieved. We fully expect the Rankine-Hugoniot relations will fail in traditional steady-state MHD models as well, since those basically capture the same transient dynamics of each field line.  While analytic treatments might suffer from erroneous assumptions about the shocks, we expect time-dependent MHD solutions to reveal the large density enhancements here predicted.  There have been several two-dimensional simulations of reconnection across a current sheet which included temperature-dependent, field-aligned thermal conduction necessary to our model \\citep{Yokoyama1997,Yokoyama1998,Chen1999b,Chen1999}.  While none included a guide field component, the reconnection outflows consisted of heat fronts ahead of isothermal slow shocks --- the same structure predicted by the shock tube model.  The simulations used fairly large initial pressures ($\\beta_0\\sim0.1$) and thus achieved rather modest shocks:  \\citet{Chen1999} observed a fifteen-fold pressure jump, which would be produced by a shock tube inflow at $M_i=2.40$.  That simulation exhibited a five-fold density enhancement across the ``isothermal shock'', twice the maximum achievable by a switch-off sock in steady state but exactly matching shock tube prediction \\eq\\ (\\ref{eq:rho_ad}).  Since the anti-parallel reconnection ($\\Delta\\theta=180^{\\circ}$) simulated by \\citet{Chen1999} falls outside the validity of the shock-tube model, this agreement must be regarded as fortuitous.  A more direct treatment of the full problem, including a guide-field, using the full set of MHD equations could bypass the limitations of the thin flux tube model.  This model requires that the plasma-$\\beta$ be small both before and after reconnection.  For negligible initial values ($\\beta_0\\ll1$), the post shock value approaches the limit, $\\beta_2\\simeq (16/3)\\sin^4(\\Delta\\theta/4)$, \\citep{Longcope2009} which exceeds $0.75$ for $\\Delta\\theta>150^{\\circ}$.  Thus reconnection at more acute angles, where $M_i$ could be large, cannot be treated using our thin flux tube model.  Indeed, there is observational evidence \\citep{Caspi2010} that many SH-LTSs achieve a plasma pressure limited by the magnetic pressure ($\\beta_2\\simeq1$).  We therefore hope in the future to use full MHD simulations to explore this reconnection regime.  In both steady or transient scenarios, the total emission measure of the hot reconnection outflow will be proportional to the mean reconnection rate $\\dot{\\Phi}$ according to \\eq\\ (\\ref{eq:EMphi_dot}).  In traditional steady-state models this is equivalent to an electric field along the X-line \\citep{Petschek1964,Vasyliunas1975,Soward1982}, directly related to reconnection microphysics.  This contrasts with patchy reconnection scenarios \\citep{Klimchuk1996,McKenzie1999} where the local electric field determines the time required to reconnect each patch, but $\\dot{\\Phi}$ is proportional to the rate patches are produced, an unrelated quantity.  While it is important to resolve this ambiguity if we are to fully understand solar flares, both scenarios lead to the same relation between $\\dot{\\Phi}$ and the flare's $EM$ for physical reasons outlined above. Moreover, the mean reconnection rate $\\dot{\\Phi}$ can be directly measured using observation of flare ribbon motion \\citep{Forbes1984,Fletcher2001,Qiu2002}. The values typically obtained are consistent with the $EM$ of typical SH-LTSs.  Such measurements have been interpreted as X-line electric fields by applying the two-dimensional steady state model \\citep{Poletto1986}.  Doing so neglects the complexities of flare ribbons and magnetic fields, which seem to suggest a less ordered, more patchy reconnection process \\citep{Fletcher2004,Longcope2007,Qiu2009}.  The constant of proportionality, $\\delta EM/\\delta\\dot{\\Phi}$, can be derived from global, observable quantities.  It consists of a factor, given in \\eq\\ (\\ref{eq:EM_fctr}), depending on pre-flare conditions and the vertical extent of the pre-flare current sheet, multiplied by a dimensionless factor depending primarily on the angle $\\Delta\\theta$  between reconnecting field lines which determines the inflow Mach number $M_i$ according to \\eq\\ (\\ref{eq:Mi}). This simple dependence can be illustrated using the $40$ MK SH-LTS studied by, \\citet{Longcope2010}, which had a constant $\\delta EM/\\delta\\dot{\\Phi}\\simeq3\\times 10^{29}$: it had an emission measure $EM\\simeq3\\times 10^{48}\\,{\\rm cm}^{-3}$ during reconnection at a mean rate $\\dot{\\Phi}\\simeq10^{19}$ Mx/sec, measured using ribbon motion. Its thirteen-fold temperature increase, from $T=3$ MK to $T=40$ MK, requires a shock of $M_i\\simeq4.7$, as would occur for reconnection between fields differing by $\\Delta\\theta\\simeq90^{\\circ}$ \\citep[roughly consistent with the pre-flare magnetic model of][]{Longcope2010}. Reading from \\fig\\ \\ref{fig:smry}, this should produce an $EM\\sim0.3$ times the factor in \\eq\\ (\\ref{eq:EM_fctr}), which must therefore be $10^{30}\\,{\\rm sec\\,Mx^{-1}\\,cm^{-3}}$.  A current sheet $h=30$ Mm high with guide field $B_g=200$ G \\citep[similar to those of][]{Longcope2010} would require a pre-reconnection density (density inside the current sheet on the un-reconnected field lines) $n_{e0}=2\\times10^{10}\\,{\\rm cm}^{-3}$, to produce the observed LTS. This is lower than the value quoted by \\citet{Longcope2010} due primarily to density enhancement above the Rankine-Hugoniot value.   According to \\fig\\ \\ref{fig:smry} the loop-top density in the SH source would be $n_e\\simeq3\\times10^{11}\\,{\\rm cm}^{-3}$, five times greater than the maximum permitted by Rankine-Hugoniot relations.  Perhaps the greatest puzzle raised by the collisional model explored here is how it might relate to flares (or flare phases) with significant non-thermal populations.  Super-hot sources often appear after the tell-tale power-law HXR spectrum has vanished \\citep{Alexander1997,Jiang2006,Longcope2010}, so the collisional and collisionless processes may occur separately.  It still seems reasonable to assume the same underlying mechanism is responsible for energy transfer in both cases.  Here we have shown how reconnection can thermalize a significant fraction of stored magnetic energy through slow magnetosonic shocks, assuming sufficiently high collision rates.  Could a related process, at lower densities, produce the non-thermal populations observed at other times and in other flares?  Except in cases of anti-parallel reconnection the slow shock is essentially a {\\em parallel shock}, whose collisionless manifestation is still poorly understood \\citep[see][for overviews of the topic and different possible approaches]{Parker1961,Sagdeev1966,Stone1985,Quest1988,Khabibrakhmanov1993}. This seems to be an avenue worth pursuing given that the present model outlines a complete energetic chain from current sheet to post-flare loops.  In any event, the large densities within the SH-LTS are sufficient to thermalize the electron population, thereby explaining the absence of a non-thermal population when they are present.  A troubling discrepancy occurs in trying to explain how bulk kinetic energy, produced at the rotational discontinuities and carried almost entirely by ions, can be transformed to electron thermal energy.  Even at the high densities observed the classical rate of ion-electron collision is too low for this transfer to be effective \\citep{Longcope2010b}.  Since all measurements available are of {\\em electron} temperatures in flares, they must be heated somehow, in spite of this theoretical hurdle.  This would not be the first instance where actual collision rates greatly exceeded those of classical Coulomb interactions.  Indeed, the resolution of this paradox may also provide a clue to the production of non-thermal electrons in cases of lower density."
    },
    "1107/1107.1748_arXiv.txt": {
        "abstract": "{  We present the discovery of four new long-period planets within the HARPS high-precision sample: \\object{HD137388}b ($M\\sin{i}$ = 0.22 $M_J$), \\object{HD204941}b ($M\\sin{i}$ = 0.27 $M_J$), \\object{HD7199}b ($M\\sin{i}$ = 0.29 $M_J$), \\object{HD7449}b ($M\\sin{i}$ = 1.04 $M_J$). A long-period companion, probably a second planet, is also found orbiting HD7449. Planets around HD137388, HD204941, and HD7199 have rather low eccentricities (less than 0.4) relative to the 0.82 eccentricity of HD7449b. {All these planets were discovered even though their hosting stars have clear signs of activity. Solar-like magnetic cycles, characterized by long-term activity variations, can be seen for HD137388, HD204941 and HD7199, whereas the measurements of HD7449 reveal a short-term activity variation, most probably induced by magnetic features on the stellar surface. We confirm that magnetic cycles induce a long-term radial velocity variation and propose a method to reduce considerably the associated noise.} The procedure consists of fitting the activity index and applying the same solution to the radial velocities because a linear correlation between the activity index and the radial velocity is found. Tested on HD137388, HD204941, and HD7199, this correction reduces considerably the stellar noise induced by magnetic cycles and allows us to derive precisely the orbital parameters of planetary companions. ",
        "introduction": "\\label{sect:1}   {The radial velocity (RV) method has been the most successful technique until now in discovering new planets\\footnote{see The Extrasolar Planets Encyclopaedia, http://exoplanet.eu}. The most effective RV machine at present, the HARPS spectrograph, now reaches the sub-meter-per-second precision level \\citep[][]{Mayor-2009b,Lovis-2006}. Although the instrumental precision is excellent, the technique has some problems in identifying very low mass planets far away from their stars because we begin then to be sensitive to the noise caused by the stars themselves. A good understanding of the stellar noise is required if we wish to a) find Earth-like planets in the habitable regions of solar-type stars using the RV technique and b) confirm the 1200 planet candidates announced by the Kepler mission team \\citep[][except for some systems showing transit timing variations]{Borucki-2011}.}  The RV technique is an indirect method that does not allow us to directly detect a planet: it measures the (gravitational) stellar wobble induced by the planetary companion orbiting its host star. Thus, the technique is sensitive not only to eventual companions, but also to stellar noise. At the sub-meter-per-second precision level, we begin to be affected by oscillation, granulation, and activity noises \\citep[][]{Dumusque-2011a, Dumusque-2011b, Boisse-2011,Arentoft-2008,Huelamo-2008,Santos-2004a, Queloz-2001}. In this paper, we are interested in these different types of activity noise.  It has already been observed and demonstrated that magnetic features such as spots or plages on the stellar surface induce a RV signature when the star is rotating. Owing to rotation, one part of the star will be blueshifted whereas the other part will be redshifted. This effect is symmetric and the average RV will be zero. Active regions, which have a different flux than the solar surface, will break this balance and introduce a variation in the RVs coupled with rotation {\\citep[][]{Boisse-2011,Queloz-2009,Henry-2002,Queloz-2001,Saar-1998,Saar-1997,Baliunas-1983,Vaughan-1981}}. We note that since this noise is proportional to the rotational velocity of the star ($v\\sin{i}$), for similar magnetic features, the RV amplitude of the effect will be more significant for rapid rotators. Active regions contain spots, which are cooler than the average stellar surface, as well as plages, which are hotter. Therefore, the noise induced by active regions will usually be compensate for, but not entirely because of variation in the surface ratio of spots to plages \\citep[e.g.][]{Chapman-2001}. The RV rms variation caused by spots and plages at the maximum of the Sun's activity is estimated to 48 cm\\,s$^{-1}$ and 44 cm\\,s$^{-1}$, respectively \\citep[][]{Meunier-2010a}. We find a similar value for spots, 51 cm\\,s$^{-1}$, when simulating the Sun's activity \\citep[][]{Dumusque-2011b}. The total effect of spots plus plages reported by \\citet{Meunier-2010a} for the same activity level is 42 cm\\,s$^{-1}$, which shows the compensation effects between the two types of active regions. This noise is induced by activity coupled with rotation and therefore appears on short-period timescales. Below we refer to this as short-term activity noise.  {Concerning magnetic cycles, the idea that long-period stellar activity variation can perturb RVs dates back to \\citet{Campbell-1988} and \\citet{Dravins-1985}. Variation in the activity of thousand of FGK dwarfs, using the Mount-Wilson S-index and the log(R'$_{HK}$) index have been studied for twenty years \\citep[][]{Hall-2007,Baliunas-1995,Duncan-1991}. However, no straightforward correlation between the long-term activity variation and the RV has been indentified until now \\citep[][]{Santos-2010b,Isaacson-2010,Wright-2005,Wright-2004,Paulson-2004,Paulson-2002}. In this paper, we show that this correlation exists and the previous non-detections can easily be explained by the precision needed and the long-term RV follow up required to detect the effect  of magnetic cycles on RVs. According to \\citet{Dravins-1985} and \\citet{Meunier-2010a}, the inhibition of convection in active regions could be responsible for the correlation.} Because solar-like stars possess an external convective envelope, granulation will appear on the surface. Hot granules of plasma come to the surface, whereas cooled plasma flows into the interior, surrounding the hot granules and forming B\\'enard cell patterns. Hot granules occupy a larger area than downward flows, which are less luminous. Because of the brightness-velocity correlation of the solar granulation \\citep[][]{Dravins-1986,Kaisig-1982,Beckers-1978}, the surface will appear blueshifted in the presence of granulation. In active regions, the convection is greatly reduced \\citep[][]{Meunier-2010a,Gray-1992,Brandt-1990,Livingston-1982,Dravins-1982} leading to a redshift of the spectrum relative to the average solar spectrum. Solar-like magnetic cycles are characterized by an increasing filling factor of magnetic regions when the activity level rises. {The total star convection is therefore lower and the star will appear redder (positive velocity) during its high-activity phase}. A positive correlation between the RVs and the activity level is then suspected \\citep[][]{Dumusque-2010,Meunier-2010a}. {This explanation is the most commonly given, although some other physical processes could be responsible for this long-term RV-activity correlation, for example the variation in surface flows proposed by \\citet{Makarov-2010}.}  We present here our detection of four new planet systems that have host stars with clear signs of activity. Three stars display solar-like magnetic cycles (long-term activity), whereas one star has short-term activity noise, related to magnetic features on its surface. In Sects. \\ref{sect:2} and \\ref{sect:3}, informations about the observations made and the stellar parameters of the hosting stars are given, and in Sect. \\ref{sect:5}, planet properties for each star are derived. Finally Sect. \\ref{sect:6} discusses our results. In addition, Table 4, 5, 6, and 7, only available in the online version and at the CDS, contains the RVs and the activity index for each star. ",
        "conclusions": "\\label{sect:6}  We have presented the detections and the orbital parameters of four planets and a companion of unknown type in four different systems. Three of these planets, HD137388 (67 M$_{\\oplus}$, 331 days), HD204941b (85 M$_{\\oplus}$, 1733 days), and HD7199b (92 M$_{\\oplus}$, 615 days) have low masses and long orbital periods in comparison with known planets\\footnote{See The Extrasolar Planets Encyclopaedia, http://exoplanet.eu}. Only three other planets have an orbital period longer than 300 days and a mass lower than 100 M$_{\\oplus}$ : HD10180\\,g \\citep[21 M$_{\\oplus}$, 601 days,][]{Lovis-2010} , HD85390\\,b \\citep[42 M$_{\\oplus}$, 788 days,][]{Mordasini-2011}, and HD10180\\,h (64 M$_{\\oplus}$, 2222 days). We note that all these planets have been found using HARPS, thus this instrument is exploring a new domain of sub-Saturn planets beyond $\\sim$\\,1\\,AU. The fourth planet, HD7449b, has a high eccentricity of 0.82 and complements the five already known planets with a higher eccentricity. Although the eccentricity is very high, the transiting probability is low because of the large separation between the star and the periastron of the planet orbit ($\\sim$\\,0.6\\,\\% for the transit and 1\\,\\% for the anti-transit\\footnote{The angle of the periastron, $\\omega$, has been taken into account in the probability computation \\citep[][]{Kane-2008}}). The discovery of these four new planets is very important to constrain and test models of planetary formation because they occupy regions of the mass-period-eccentricity space where only a few other discoveries have been made until now.  We found a second companion orbiting HD7449. Given that the data only cover a small fraction of the orbital period, the nature of this object, planet, brown dwarf, or small stellar companion is hard to determine. A Monte Carlo simulation shows, however, that a planetary companion with a mass similar to 2\\,$M_J$ and a period close to 4000 days is the most probable. {Given the high eccentricity of the inner planet, it is very likely that a strong interaction occurs between the two companions, although much more data are needed to define properly this interesting multi-planetery system.}   The planets announced in this paper for the first time have been discovered even though their host stars display clear signs of activity. We have found that HD7449 exhibits signs of short-term activity, whereas HD7199, HD137388, and HD204941 have solar-like magnetic cycles.  When we examined the periodogram of HD7449 RV residuals we are able to identify two significant peaks close to a 14 day period, which is the estimated rotational period of the star, 13.30 $\\pm$ 2.56 days. A similar signal appears in the periodogram of the activity index, although does not appear to be very significant. In addition, no signal was detected in either the BIS span or the FWHM. Therefore, a low-mass planet might be present. Looking at different chunks to see whether the 14-day period signal fitted stays in phase or not, we have arrived at the conclusion that the main signal is most probably caused by short-term activity and not induced by a planet on a circular orbit. However, a low-mass eccentric planet might still be present at this period and more measurements are needed to clearly define the nature of this 14-day signal.  When examining the RVs and the fitted planets for HD7199, HD137388, and HD204941, it is clear that magnetic cycles induce RV variations that could be misinterpreted as long-period planetary signature. Therefore, the long-term variations in the activity index have to be studied properly to distinguish between the real signature of a planet and long-term activity noise. When the long-term activity index variation evolves smoothly, it can be readily fitted and thus effect of magnetic cycles can be removed from RVs. This has been done for HD7199, HD137388, and HD204941 and gives a good estimation of the planet parameters, with rather small errors.  {In \\citet{Meunier-2010a}, the authors argued that the Sun should show RV variations of 10\\,m\\,s$^{-1}$ over its cycle. In this paper, we have confirmed that this behavior is seen on other solar-like stars and one can ask whether all stars showing magnetic cycles have a long-term RV variation induced by the long-term activity variation. The high precision HARPS sample, composed of 451 stars, provides a good set of measurements to search for this activity-RV correlation. However, owing to irregular sampling, this property can only be studied for a small set of stars. In preliminary work \\citep[][]{Dumusque-2010}, we demonstrated that the magnetic cycle-RV correlation can be observed for 31 comprehensively observed stars and that in addition, RVs of G dwarfs are more sensitive to magnetic cycles than RVs of K dwarfs. A more complete study is in progress and will be soon published (Lovis et al. in prep.).}"
    },
    "1107/1107.1054_arXiv.txt": {
        "abstract": "We present here the results of the observation of CTB 37A obtained with the X-ray Imaging Spectrometer onboard the {\\it Suzaku} satellite. The X-ray spectrum of CTB 37A is well fitted by two components, a single-temperature ionization equilibrium component (VMEKAL) with solar abundances, an electron temperature of $kT_{\\rm e}\\sim0.6$ keV, absorbing column density of $N_{\\rm H}\\sim3\\times10^{22}$ ${\\rm cm^{-2}}$ and a power-law component with photon index of $\\Gamma$ $\\sim 1.6$. The X-ray spectrum of CTB 37A is characterized by clearly detected K-shell emission lines of Mg, Si, S, and Ar. The plasma with solar abundances supports the idea that the X-ray emission originates from the shocked interstellar material. The ambient gas density, and age of the remnant are estimated to be $\\sim$1$f^{-1/2}$${\\rm cm^{-3}}$ and $\\sim$3$\\times10^{4}f^{1/2}$ yr, respectively. The center-filling X-ray emission surrounded by a shell-like radio structure and other X-ray properties indicate that this remnant would be a new member of mixed-morphology supernova remnant class. ",
        "introduction": "CTB 37A (also called G348.5+0.1, $\\rmn{RA}(2000)=17^{\\rmn{h}} 14^{\\rmn{m}} 06^{\\rmn{s}}$, $\\rmn{Dec.}~(2000)=-38\\degr 32\\arcmin$) was discovered by \\citet{b5} in the radio band and has a shell-type morphology with an angular size of 15 arcmin. From Very Large Array observations at wavelengths of 6, 20, and 90 cm, \\citet{b6} reported that this supernova remnant (SNR) is expanding in an inhomogeneous region, and is part of a complex composed of three SNRs: CTB 37A, G348.7+0.3 (also called CTB 37B), and G348.5-0.0. From the {\\it ASCA} Galactic plane survey data of G348.5+0.1, \\citet{b28} found that the X-ray spectra of the SNR was heavily absorbed by interstellar matter with $N_{\\rm H}\\sim2\\times10^{22}$ ${\\rm cm^{-2}}$ and the size of X-ray emission was comparable to its radio structure. \\citet{b31} detected OH masers with velocities at about 20 and 60 km s$^{-1}$, in the direction of CTB 37A at 1720 MHz. \\citet{b53} surveyed the environment of this remnant with associated OH 1720 MHz masers in the CO J=1-0 transition with the 12 Meter Telescope of the NRAO. They reported that a number of molecular clouds are interacting with the SNR shock fronts. \\citet{b20} using {\\it Chandra} and {\\it XMM-Newton} data showed the presence of thermal X-rays from the Northeast part, an extended non-thermal X-ray source, CXOU J171419.8-383023 in the Northwest part and a ${\\gamma}$-ray source, HESS J1714-385, coincident with the remnant. They found a high absorbing column density of $N_{\\rm H}\\sim3\\times10^{22}$ ${\\rm cm^{-2}}$. They claimed that the observed X-ray morphology was a result of interaction with the inhomogeneous medium surrounding the remnant and this inhomogeneity was also responsible for the break-out radio morphology. \\citet{b51} using observations with the Fermi-LAT, have revealed ${\\gamma}$-ray emission from this SNR; the spectrum of the source coincident with CTB 37A was fitted by a power-law (PL) model with an exponential cutoff at energy $E_{\\rm cut}$=4.2 GeV. Considering the lack of evidence for contribution by a pulsar and the presence of maser emission for the remnant they proposed that ${\\gamma}$-rays result from SNR-molecular clouds interactions.  The distance to the CTB 37A has been estimated from 21 cm absorbtion measurement to be in between 6.7-13.7 kpc by \\citet{b52}. From velocity measurements of molecular clouds associated with the remnant, \\citet{b53} adopted a distance of 11.3 kpc. So we will use d=11.3 kpc for our calculations throughout this work.  The location of CTB 37A containing OH maser sources \\citep{b31}, 1FGL J1714.5-3830 \\citep{b51}, X-ray source CXOU J171419.8-383023, and HESS J1714-385 \\citep{b20}, as well as two additional SNRs (G348.5-0.0 and G348.7+0.3) make this SNR interesting and important. In addition, its morphology implies a similarity to the recently proposed group of mixed-morphology (MM) SNRs, increasing its importance. High quality imaging and spectra obtained from the data provided by X-ray observatory {\\it Suzaku} are used to produce the results in this work.  The organization of this paper is as follows: we describe the {\\it Suzaku} X-ray observation of CTB 37A, including details of data reduction in Section 2. The image and spectral analysis are given in Sections 3 and 4, respectively. Finally, considering its morphology and investigating radial variation of the electron temperature, we discuss the implications of MM class, the nature of thermal and non-thermal components of CTB 37A in Section 5.  \\section[]{Observation and Data reduction}  {\\it Suzaku} \\citep{b17} is the fifth Japanese X-ray astronomy satellite launched on 2005 July 10. {\\it Suzaku} has observed CTB 37A on 2010 February 20. The observation ID and exposure time are 504097010 and 53.8 ks, respectively. The X-ray Imaging spectrometer (XIS; \\citet{b12}) consists of four sets of X-ray CCD camera system (XIS0, 1, 2 and 3). XIS0, 2, and 3 have front-illuminated (FI) sensor and provide coverage over the energy range 0.4$-$12 keV, while XIS1 has a back-illuminated (BI) sensor providing greater sensitivity at lower energies (0.2$-$12 keV). The XIS has a field of view (FOV) of 17.8$\\times$17.8 arcmin$^{2}$ ($1024\\times1024$ pixels). Each XIS CCD has an $^{55}$Fe calibration source, which can be used to calibrate the gain and test the spectral resolution of data taken using this instrument. The XIS2 sensor was available only until 2006, therefore, we use data of XIS0, XIS1, XIS3.  Reduction and analysis of the data were performed following the standard procedure using the {\\sc headas} software package of version 6.5 and spectral fitting was performed with {\\sc xspec} version 11.3.2 \\citep{b2}. The XIS was operated in the normal full-frame clocking mode, with the standard $3\\times3$ and $5\\times5$ editing mode. We generated XIS response matrices using the {\\sc xisrmfgen} software, which takes into account the time-variation of the energy response. As for generating ancillary response files (ARFs), we used {\\sc xissimarfgen} \\citep{b7}. The latest version of the relevant {\\it Suzaku} CALDB files were also used. ",
        "conclusions": "In this work, we provide a description of the X-ray emission of CTB 37A based on {\\it Suzaku} archival data. We obtained a clear image and high quality spectra of diffuse X-ray emission. We have examined the thermal and non-thermal emissions coming from the remnant. The X-ray spectrum of CTB 37A is characterized by thermal emission dominated by K-shell emission lines of Mg, Si, S, and Ar which are clearly detected.  As seen from Fig. 2, the radio emission of CTB 37A comprises a partial shell towards the north and east and an extended outbreak to the south, while the X-ray emission has a deformation along the southwest limb, where the morphology appears indented. The reason for such a deformation may be due to the inhomogeneous medium along this specific region and this may well be supported by the fact that remnant is close to the Galactic plane and several OH masers at 1720 MHz are detected towards CTB 37A \\citep {b31} (shown by crosses). It may also be relevant that $\\gamma$-ray emission thought to be associated with the interaction between CTB 37A and dense surrounding material that has been detected with the {\\it Fermi LAT} \\citep {b51}.  \\subsection{Implications for mixed-morphology}  SNRs were originally divided into shell-like, Crab-like (plerionic), and composite (shell-like containing plerions) remnants \\citep {b42} according to their X-ray morphology. Recently, an additional MM class (also called thermal composite) appeared, which are center-filled in X-rays and shell-like at radio wavelengths \\citep {b23, b24, b34}. As seen from Fig. 2, CTB 37A has a shell-like morphology in the radio band while centrally-filled in X-ray band, in this regard, our first impression is CTB 37A seems to be a MM SNR. Examples of well-known MM SNRs include W28 \\citep {b48}, G290.1-0.8 \\citep {b49}, and IC 443 \\citep {b4}. \\citet {b3} reported important results about this class by compiling a list of 26 MM SNRs.  The X-ray characteristics of MM remnants have been defined by \\citet {b34} as follows: (1) The radial temperature distribution is relatively flat; (2) The X-ray emission arises primarily from shocked interstellar material, and not from ejecta; (3) The remnants are typically located close to molecular clouds or very dense regions; and (4) The dominant X-ray emission is thermal in nature. Subsequent studies indicated that MM SNRs had a complex plasma structure with multiple components (e.g. \\citet {b48}) and enhanced abundances (e.g. \\citet {b50,b49}), and they have evolved over $\\sim$$10^{4}$ yr, which means that the plasma is in CIE or an overionization condition \\citep {b47}. As can be seen from Fig. 4, the annular analysis of CTB 37A may well be indicating a very small scale radial variation in its temperature between the selected regions. The best-fitting metal abundances are found to be solar in general, confirming  the absence of ejecta contamination in selected regions. It may support the idea that the X-ray emission originates from the shocked interstellar material. The plasma of CTB 37A is in a collisional ionization equilibrium condition and is located in a region with density variation, possibly associated with molecular clouds. The plasma has thermal and non-thermal emission, but the emission is dominated by thermal component ($\\sim90$ per cent of the total X-ray flux). These X-ray properties of CTB 37A exemplify the typical characteristic of MM SNRs as defined by \\citep {b34}.  There are a few models that can produce centrally enhanced X-ray emission such as  evaporation of clouds left relatively unspoiled after the passage of the SNR blast wave (e.g. \\citet {b9}) and ``fossil'' thermal radiation that is detectable as thermal X-rays from the hotter interior as the shell of an expanding  SNR cools below $\\sim$$10^{6}$ K and  becomes invisible due to interstellar absorbtion (e.g. \\citet {b42}), as the SNR evolves, the temperature and density of the hot interior plasma gradually become uniform through thermal conduction (e.g. \\citet {b10}) or evolution in a medium with a density gradient viewed along the line of sight \\citep {b26}. The evaporation model requires dense clouds, the thermal conduction model requires a relatively high density ambient medium. CTB 37A is located in a region of greatly varying density, with OH maser sources indicating interaction with molecular clouds. In this regard, the center-filled X-ray morphology of CTB 37A is consistent with evaporating clouds model. The radial temperature variation in the plasma of CTB 37A ($kT_{\\rm e}$ $\\sim$ 0.6-0.8 keV) is consistent with other MM SNRs such as W44 \\citep {b29,b21}, 3C391 \\citep {b32,b33} and HB21 \\citep {b27}. Very small temperature variation in the plasma of CTB 37A as shown Fig. 4 can be explained by both evaporation and thermal conduction models. Future deep X-ray observations and detailed spectral analysis of this remnant would give more detailed information to compare with theoretical models that produce MM SNRs.  \\subsection{Thermal component}  The X-ray emission of CTB 37A is dominated by thermal emission that can be best described by an absorbed CIE plasma model (VMEKAL) with an absorbing column density of $N_{\\rm H}\\sim 3\\times10^{22}$ $\\rm cm^{-2}$, an electron temperature of $kT_{\\rm e}\\sim0.6$ keV, and solar abundances of Mg, Si, S, and Ar, which indicate a shocked interstellar/circumstellar material origin.  For full ionization equilibrium, the ionization timescale, $\\tau=n_{\\rm e}t$, is required to be $\\geq$ $10^{12}$ $\\rm cm^{-3}$s, where t is the plasma age or the time since the gas was shock-heated \\citep {b14}. To determine the age of the remnant, $n_{\\rm e}$ should be estimated from the emission measure, $n_{\\rm e}n_{\\rm H}V$, which is related to the normalization of the VMEKAL model according to the equation, norm=$n_{\\rm e}n_{\\rm H} V$/($4\\pi d^{2}$$10^{14}$), where V is the X-ray emitting volume, $n_{\\rm H}$ is the volume density of hydrogen and d is the distance. For simplicity, we assumed the emitting region to be a sphere of radius 5.5 arcmin. Considering the possibility that less than the entire volume is filled, we write the volume V=$V_{\\rm s}f$, where $V_{\\rm s}$ is the full spherical volume, $f$ is the filling factor. We then carry the $f$ factor through our calculations to show the explicit dependence of each derived quantity on this factor. Knowing that the SNR is at a distance of 11.3 kpc and $n_{\\rm e}=1.2n_{\\rm H}$, we estimated the emission volume to be $V\\sim 6.7\\times10^{59}f$ ${\\rm cm^{3}}$. Consequently, we find an ambient gas density of $\\sim$1$f^{-1/2}$ ${\\rm cm}^{-3}$ and age of $\\sim$$3\\times10^{4}f^{1/2}$ yr (assuming $n_{\\rm e}t\\sim 1\\times10^{12}$ $\\rm cm^{-3}$s), implying that CTB 37A is a middle-aged SNR. Finally, we calculated total mass of the X-ray emitting plasma, $M_{\\rm x}$, by $M_{\\rm x}$=$m_{\\rm H}n_{\\rm e}V$$\\sim 530 f^{1/2}{M\\sun}$, where $m_{\\rm H}$ is mass of a hydrogen atom, $\\mu$=0.604 is the mean atomic weight.  \\subsection{Power-law component}  The Suzaku X-ray spectral data of CTB 37A is well fitted with a thermal component and an additional hard component. There could be a few reasons for the hard X-ray emision: (i) an association with a classical young pulsar, (ii) a contribution from an extended non-thermal X-ray source (CXOU J171419.8-383023), (iii) overionization of the plasma which produces excess hard emission as has been the case for IC443 \\citep {b60}. The hard component is well fitted by a PL model with a photon index value of $\\sim$1.6. This value is consistent with that of classical young pulsar value ranging in between 1.1 and 1.7 \\citep {b37}. However, there is no pulsar reported that is associated with this remnant. In the Northwest region of the CTB 37A, an extended non-thermal X-ray source (CXOU J171419.8-383023) is reported ($\\rmn{RA}(2000)=17^{\\rmn{h}} 14^{\\rmn{m}} 20^{\\rmn{s}}$, $\\rmn{Dec.}~(2000)=-38\\degr 30\\arcmin 20\\arcsec$) by \\citet {b20}. In their work, a non-thermal emission from the source  with a spectral index of $\\sim$1.32  is found which is lower (harder) than our best-fitting value of $\\sim$1.6. Although we excluded CXOU J171419.8-383023 (with radius $2.1$ arcmin) from our spectra during our spectral analysis, our fits required a non-thermal component. To investigate this, we performed spectral analysis also for individual regions by selecting small rectangular regions that are being further away from the known extended non-thermal source. The non-thermal flux is found to be stronger for the selected small regions nearby the source compared to the ones further away. We have obtained an unabsorbed flux value of $F_{\\rm x}\\sim1.4\\times10^{-9}$ erg $\\rm s^{-1}$$\\rm cm^{-2}$ for the non-thermal extended source in  the 0.3$-$10 keV energy range. When we compare it with the unabsorbed flux value ($F_{\\rm x}\\sim0.14\\times10^{-9}$ erg $\\rm s^{-1}$$\\rm cm^{-2}$) of PL component of best-fitting, we find a factor of ten difference between them. This difference indicates that the extended source is most likely the origin of the PL component and emission scattered from the source into the field of the rest of the remnant by a broad point spread function of the {\\it Suzaku} mirrors.  The spectral studies of CTB 37A indicate that the plasma is best described by a thermal component in CIE condition with solar elemental abundances and a non-thermal component with a photon index of $\\sim$1.6. Thermal emission possibly originates  from the shocked interstellar material with ambient gas density of $\\sim$1$f^{-1/2}$${\\rm cm^{-3}}$. The best spectral fits  require an ionization timescale  of  $\\tau\\geq$ $10^{12}$ $\\rm cm^{-3}$s, implying an age of $\\sim$3$\\times10^{4}f^{1/2}$ yr. The origin of the power-law component is more likely the effect of the contribution from the extended source (CXOU J171419.8-383023) located  in the Northwest part of the remnant. CTB 37A is most likely a new member of mixed-morphology SNR."
    },
    "1107/1107.5488_arXiv.txt": {
        "abstract": "\\sgr\\ is a transient Soft Gamma-ray Repeater which underwent a major outburst in June 2009, during which the emission of short bursts was observed. Its properties appeared quite typical of other sources of the same class until long-term X-ray monitoring failed to detect any period derivative. The present upper limit on $\\dot P$ implies that the surface dipole field is $B_p\\lesssim 7.5\\times 10^{12}\\ {\\rm G}$ \\cite[][]{rea10}, well below those measured in other Soft Gamma-ray Repeaters (SGRs) and in the Anomalous X-ray Pulsars (AXPs), a group of similar sources. Both SGRs and AXPs are currently believed to be powered by ultra-magnetized neutron stars (magnetars, $B_p\\approx 10^{14}$--$10^{15}\\ {\\rm G}$). \\sgr\\ hardly seems to fit in such a picture. We show that the magneto-rotational properties of \\sgr\\ can be reproduced if this is an aged magnetar, $\\approx 1\\ {\\rm Myr}$ old, which experienced substantial field decay. The large initial toroidal component of the internal field required to match the observed properties of \\sgr\\ ensures that crustal fractures, and hence bursting activity, can still occur at present time. The thermal spectrum observed during the outburst decay is compatible with the predictions of a resonant compton scattering model (as in other SGRs/AXPs) if the field is low and the magnetospheric twist moderate. ",
        "introduction": "Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs) form a small, but rapidly expanding, class of isolated neutron star (NS) sources characterized by the emission of short ($\\approx 0.1$~s), energetic ($\\approx 10^{40}$ \\ergs) bursts of X-rays. Their persistent luminosity ($L_X\\approx 10^{32}$--$10^{36}$ \\ergs\\ in the $\\sim 0.5$--10 keV range) is often variable, with flux enhancements up to several hundreds during the outburst phases of transient sources \\cite[e.g.][]{reaesp11}. SGRs and AXPs share similar timing properties, with periods in a narrow range ($P\\sim 2$--12~s), and large period derivatives ($\\dot{P} \\approx 10^{-13}$--$10^{-10} \\ {\\rm s \\ s}^{-1}$). In most of these sources the X-ray luminosity exceeds the rate of rotational energy losses ($\\dot E\\approx 10^{32}$--$10^{35}$ \\ergs), and highly variable radio activity has been detected in three objects. X-ray spectra are often characterized by a thermal (BB; $kT\\approx 0.5$ keV) plus a high-energy power-law (PL; photon index $\\Gamma\\approx 1.5$--3) component \\cite[e.g.][]{rea08}, or by two thermal components \\cite[e.g.][]{gotthalp07,tiengo08}. Both the spectral parameters and pulse profiles appear to vary with the source flux \\cite[see e.g.][for reviews on SGRs/AXPs properties]{woodsthomp06, mereghetti08, reaesp11}\\footnote{see also {\\tt http://www.physics.mcgill.ca/\\textasciitilde pulsar/magnetar/\\\\main.html} for an updated catalogue of SGRs/AXPs}.  The large values of the magnetic field derived from the standard dipole formula ($5\\times 10^{13}\\ {\\rm G}\\lesssim B_p \\lesssim 10^{15}\\ {\\rm G}$), the lack of detected stellar companions and their large X-ray output as compared to $\\dot E$ led to the suggestion that SGRs and AXPs are powered by a young (age $\\approx 10^3$--$10^4$ yr), ultra-magnetized neutron star, or magnetar \\cite[][]{duncan92, thompson93}. Indeed, the magnetar scenario has been largely successful in explaining many of the observed properties of SGRs/AXPs, including those of the bursts \\citep{td95} and of the persistent emission \\cite[the twisted magnetosphere model,][and references therein]{thompson02, zrtn09, alba10}.  Despite SGRs and AXPs are far from being a homogeneous class, in particular the inferred surface dipolar field spans nearly two orders of magnitude, their observational behaviour is now commonly associated to that of (active) magnetars, to the point that often the terms SGR/AXP and magnetar are used as synonyms. This, actually, reflects the original definition of a magnetar as a NS which is powered by its (large) magnetic field \\cite[][]{thompson93}. In this respect, a super-strong magnetic field is not per se a sufficient condition for triggering SGR/AXP-like activity, as testified by the existence of NS sources, for instance most of the so-called High-B radio pulsars \\cite[HBPSRs; e.g.][]{kaspi10}, and possibly some of the thermally emitting isolated NSs \\cite[XDINSs; e.g.][]{turolla09}, with surface magnetic fields comparable to those of SGRs/AXPs but having substantially different properties and not showing any bursting/outbursting activity over the $\\sim 10$--20 yr time span during which they were observed.  Over the last few years, the family of the magnetar candidates grew with the addition of several new objects, most of them transient. X-ray bursting activity or peculiar radio emission similar to that of SGRs/AXPs was detected in allegedly rotation-powered pulsars, as PSR J1846-0258 \\cite[][]{gav08, kum08} and PSR 1622-4950 \\cite[][]{lev10}. This suggests that magnetic energy substantially contributes to their emission at certain stages. The magnetic field inferred from their rotation properties ($B_p \\sim 0.5$--$3\\times 10^{14}\\ {\\rm G}$) is, in fact, close to that of SGRs/AXPs, contributing to the widespread belief that magnetar-like activity has to be associated to super-strong magnetic fields, typically higher than the quantum field $B_Q\\simeq 4.4\\times 10^{13} \\ {\\rm G}$.  In this respect, the recent discovery of a low-field SGR, \\sgr\\ \\cite[][]{rea10}, came as a surprise. \\sgr\\ was observed for the first time in June 2009 when it entered an outburst state during which X-ray bursts were detected \\cite[][]{vdhorst10}. The enhanced flux level allowed for a measure of the source periodicity, $P\\sim 9.1$~s, but, despite the source was then monitored in X-rays for $\\sim 500$~d, no significant evidence for a period derivative was found \\cite[][and references therein]{esp10, rea10}. The published upper limit is $\\dot P\\lesssim 6\\times 10^{-15}$ s$\\,\\rm s^{-1}$, leading to an inferred dipole field $B_p\\lesssim 7.5\\times 10^{12}$~G.  The very low magnetic field of \\sgr\\ (when compared to other SGRs/AXPs) raises a number of questions as to if, and how, the observed phenomenology of this source can be accommodated within the magnetar picture. A crucial point is whether a NS with a surface field well below $10^{13}$~G, can harbor an internal toroidal magnetic field strong enough to produce crustal displacements, which are believed to be responsible for the bursting/outbursting episodes in  SGRs/AXPs \\cite[][]{td95,thompson02, belob09}. In this paper we address this issue and discuss both the spectral and timing properties of \\sgr\\ in the framework of an aging magnetar. ",
        "conclusions": "\\sgr\\ is the first ``low-field'' soft gamma repeater/anomalous X-ray pulsar ever discovered. Even if other sources of the same class with a relatively weak $B$-field were already known \\cite[e.g. 1E 2259+586,][]{gk02}, the case of \\sgr\\ is extreme and stimulated considerable interest. With a surface dipole field $\\lesssim 7.5\\times 10^{12}\\ {\\rm G}$, \\sgr\\ seems to challenge an interpretation in terms of the ``conventional'' magnetar model, and not just because its magnetic field is below the quantum critical field, a quite irrelevant fact per se.  The present upper limit on the surface field in \\sgr\\ is based on the spin-down measure, which traces the dipole component. It is actually possible that the external magnetic field in NSs, and in SGRs/AXPs in particular, is more complex than a simple dipole. Higher order multipoles can substantially contribute to the field near the surface and, because they fall off more rapidly with radius, induce negligible spin-down. \\sgr\\ may, then, possess a much higher surface $B$-field than indicated by its spin-down rate.  A high surface field in the form of multipolar components hints toward the presence a large internal field which can produce crustal motions and bursting activity, making \\sgr\\ not dissimilar (a part from the magnetic field topology) from other magnetar sources. It does not, however, explain the rotational properties of the source. With a dipole field below $10^{13}\\ {\\rm G}$, in fact, it would be impossible to slow down the star to the observed 9.1 s period in less than $\\sim 24$ Myr, a time much longer than the estimated age of other SGRs/AXPs, unless \\sgr\\ had an exceptionally long period at birth. A possible solution was suggested by \\cite{alpar11} who have shown that a if \\sgr\\ was born with a period $> 70$ ms and a low dipolar field ($B\\approx 10^{12}\\ {\\rm G}$), the torque exerted by a fallback disk can spin down the star to the present period in $\\gtrsim 10^5$ yr, if $\\dot P$ has to be below the observed upper limit\\footnote{A somehow similar scenario in which SGRs/AXPs are NSs with a low dipole field and super-strong multipolar components powered by accretion from a fallback disk was recently proposed by \\cite{true10}.}. The survival of multipolar field components $\\approx 100$ times stronger than the dipole over such a time span may, however, be an issue in the light of the known evolution of the dipolar field in magnetars (see \\S \\ref{magevol}).  In this paper we explored a different possibility, i.e. that \\sgr\\ is an old neutron star born with a super-strong magnetic field which experienced field decay over a time $\\approx 10^6$ yr. Our results show that magneto-dipolar braking can effectively spin down the star to a period $\\sim 10$ s with $\\dot P\\lesssim 10^{-15}\\ {\\rm s\\,s}^{-1}$ provided that the initial internal toroidal field is large enough, $B_{tor}(t=0)\\gtrsim 10^{16}\\ {\\rm G}$. At the same time the initial external dipole field has to be $\\lesssim 2$--$3\\times 10^{14}\\ {\\rm G}$ in order to prevent the NS from spinning down too fast. In particular, the model with $B_{tor}(t=0)=4\\times 10^{16}\\ {\\rm G}$ and $B_{p}(t=0)=2.5\\times 10^{14}\\ {\\rm G}$ gives $B_p\\sim 2\\times 10^{12}\\ {\\rm G}$, $P\\sim 9$ s and $\\dot P\\sim 4\\times 10^{-16}\\  {\\rm s\\,s}^{-1}$ at an age of $\\sim 1.5 $ Myr, in very good agreement with the current observational picture of \\sgr. The predicted quiescent luminosity at the same age is $L_X\\sim 10^{31}$ \\ergs, again in agreement with the  current luminosity of the source, the faintest measured so far and possibly close to the quiescent value \\cite[$L_X\\sim 6\\times 10^{31}$ \\ergs\\ for a distance of 2 kpc;][]{rea10}. We note in this respect that the luminosity drops very quickly for ages $\\gtrsim 10^6$ yr (top left panel of figure \\ref{mag-rot-evol}), so the present value could well be below $\\sim 10^{31}$ \\ergs\\ if the age is only slightly above 1.5 Myr. A somehow longer age would have negligible impact on the predicted $P$, $\\dot P$ and $B_p$ which already reached a nearly constant value (see again figure \\ref{mag-rot-evol}). We also note that the measured flux in the $\\sim 0.5$--10 keV band may not be representative of the bolometric luminosity for an old object like \\sgr. If the NS surface temperature is $\\lesssim 10^6$ K, in fact, the quiescent thermal emission would be too soft (and absorbed) to be clearly detectable. It is then possible that the observed X-ray flux only constrains the ``outburst'' luminosity and the genuine quiescent emission may go undetected even if it is $\\approx 10^{31}$ \\ergs.  Given the large internal toroidal field at birth, our fiducial model for \\sgr\\ retains the capability to induce crustal fractures, and hence to produce bursts, even at quite late times. A comparison of our model with one of the cases investigated in detail by \\cite{pern11} at a comparable age ($\\sim 1.5$ Myr), indicates that the recurrence time between bursting/outbursting events for an object like \\sgr\\ is a few tens of years. Crustal displacements are accompanied by a twisting up of the magnetosphere and the formation of a high-energy spectral tail, which is observed in most SGRs/AXPs. We have shown, however, that a source with a spectral distribution consistent with a double blackbody, as follows from  the analysis of a series of X-ray observations taken during the outburst decay of \\sgr, can be modelled as well by a resonant cyclotron scattering model if the twist is moderate (twist angle $\\lesssim 0.5$ rad) and the surface field is low, $\\lesssim 5\\times 10^{12}\\ {\\rm G}$, as predicted by our evolutionary calculations.  It has been already noticed that the SGRs/AXPs which exhibit the larger flux variations seem to be those with the lower dipole fields \\cite[][]{esp09}. \\sgr\\ provides a further case for such a correlation. In the magnetar scenario the occurrence and energetics of outbursts are dictated by the internal field and the peak luminosity $L_{max}$ depends on the (maximum) energy stored locally prior to the event. $L_{max}$ is not much sensitive to the external dipole field. If the energy released in an event is roughly similar in all sources (as in the case the mechanism is gated e.g. by the crustal yield), the ratio between the peak and persistent fluxes has to be much higher in the less active, old objects than in active, young ones. This is because $L_{max}$ is similar for all of them, but the quiescent luminosity is much lower for old sources. In the latter, field decay had the time to reduce the dipole field to rather low values, so old, low-field objects (but with a sufficiently high internal field) appear as transient sources. It is intriguing that \\sgr\\ is the source with the lowest field known and at the same time a most extreme transient, with a peak-to-persistent flux ratio $\\sim 1000$."
    },
    "1107/1107.5441_arXiv.txt": {
        "abstract": "\\medskip \\noindent It is well known that the electroweak phase transition (EWPhT) in extensions of the Standard Model with one real scalar singlet can be first-order for realistic values of the Higgs mass. We revisit this scenario with the most general renormalizable scalar potential systematically identifying all regions in parameter space that develop, due to tree-level dynamics, a potential barrier at the critical temperature that is strong enough to avoid sphaleron wash-out of the baryon asymmetry. Such strong EWPhTs allow for a simple mean-field approximation and an analytic treatment of the free-energy that leads to very good theoretical control and understanding of the different mechanisms that can make the transition strong. We identify a new realization of such   mechanism, based on a flat direction developing at the critical temperature, which could operate in other models. Finally, we discuss in detail some special cases of the model performing a  numerical calculation of the one-loop free-energy that improves over the mean-field approximation and confirms the analytical expectations. ",
        "introduction": "The search for physics beyond the Standard Model (SM) has strong theoretical and experimental motivation and the simplest extension is to enhance the SM by a scalar gauge singlet degree of freedom. This minimalistic model (and its cousins with a complex singlet or supersymmetric versions of it) can be very successful in explaining various phenomena that cannot be explained by the SM: dark matter~\\cite{Bento:2000ah}-\\cite{Ponton}, spontaneous $B-L$ breaking~\\cite{Majoron}-\\cite{EWPHTMajoron} and the baryon asymmetry of the Universe~\\cite{Huber:2000mg, EWPHTnMSSM}, often leading to characteristic collider signatures~\\cite{Datta:1995qw}-\\cite{Gripaios:2009pe}.  One prominent difference between the SM and its singlet extensions is the following. While in the SM  the LEP bound on the Higgs mass ($M_h>114.4$ GeV \\cite{LEPMh}) implies that the electroweak phase transition (EWPhT) is not first-order but a smooth crossover \\cite{crossover}, the addition of a singlet can lead to strongly first-order  EWPhTs~\\cite{AH}-\\cite{Ashoorioon:2009nf} for realistic values of the scalar masses. Moreover, with such strong EWPhTs, not only the observed baryon asymmetry of the Universe can be explained through electroweak (EW) baryogenesis  (provided the model also contains additional sources of $CP$ violation) but the process of EW symmetry breaking can also leave the trace of a  stochastic background of gravitational waves~\\cite{Apreda:2001us}.  The aim of the present work is to revisit the EWPhT in the most general renormalizable extension of the SM with one additional real scalar singlet. Although this issue has been studied in the past~\\cite{AH}-\\cite{Ashoorioon:2009nf}, (both numerically and analytically at different levels of generality), we believe that a thorough  \\emph{analytical} understanding of the rich spectrum of possibilities for a strong EWPhT this model offers is still lacking in the literature. The analysis that comes closest to this task is Ref.~\\cite{Ramsey}, over which we will improve in a number of aspects.  In the SM and some extensions of it, a first-order EWPhT is caused by the thermal effects of bosons coupled to the Higgs, that generate a cubic term in the Higgs scalar potential. Although this can be successful in many cases, it requires sizeable couplings of these bosons to the Higgs and the effect can be screened by thermal masses when daisy resummation is taken into account. In this article, we concentrate on EWPhTs for which the barrier separating broken and symmetric vacua is not generated by the previous thermal cubic correction but rather by tree-level effects. These tree-level effects lead in general to stronger EWPhTs as $v_c=v(T_c)$, the Higgs vacuum expectation value (VEV) at the critical temperature $T_c$ (that controls through the celebrated ratio $v_c/T_c$ the sphaleron erasure of the baryon asymmetry), is now proportional to some $T$-independent dimensionful parameter in the potential; hence $v_c/T_c$ can become potentially very large for small $T_c$.\\footnote{Strong EWPhTs are particularly welcome if, as suggested by \\cite{DeSimone}, magnetic fields generated at the EWPhT lower the sphaleron energy so that larger values of $v_c/T_c$ are required to avoid baryon washout.} The parameter space of this model is quite rich and we will see that these tree-level barriers are not necessarily related to the presence of cubic terms in the potential, as is often assumed.  We begin in Section~\\ref{sec:PhT} by studying the structure of the tree-level scalar potential of the model. In particular we are interested in potentials where the EW breaking and preserving minima are degenerate, since this is the situation that arises for strong phase transitions. Indeed, in this case it is well justified to use the mean-field approximation for the free-energy, which will have the same structure as the tree-level potential (with temperature-dependent parameters). Differently from previous analyses, we will introduce a novel set of parameters particularly convenient for the discussion of the vacuum structure of this kind of potentials (and which might also be of use for phenomenological studies of the scalar sector of these models). In spite of the large number of free parameters (eight) we have to deal with, this new parametrization will allow us to identify very easily, and analytically, the structure of the potential: its stable minima and the existence of a barrier between them.  As a result, beside developing a better understanding of the ingredients necessary to get a strong EWPhT in this model, we will find new scenarios with strong EWPhTs that had not been identified before (involving in particular flat directions at the critical temperature).  In Section~\\ref{sec:strong_PT} we discuss thermal corrections to the scalar potential and explain our strategy to search for regions in parameter space with strong EWPhTs, which we summarize in Table~\\ref{Table1}. The idea is to start from a potential with degenerate broken and unbroken minima and a barrier between them: this is identified with a potential at $T_c$ which gives a strong EWPhT. Its parameters are then evolved to lower $T$ to find their values at $T=0$, where they can be used for phenomenological purposes. After identifying the regions in parameter space that give a strong EWPhT, we then perform a more precise analysis  including one-loop corrections (at $T=0$ and finite temperature)  without resorting to high-$T$ expansions and including daisy resummation. Although the strength of the EWPhT is somewhat reduced with respect to tree-level, one still gets sizable values for it. Our results confirm then the expectations based on the tree-level analysis.  In the rest of the paper we apply these tools to particular realizations of the model  previously considered in the literature: the $\\mathbf{Z}_2$-symmetric case in Section~\\ref{sec:Z2}; a particular supersymmetric incarnation in Section~\\ref{sec:nmssm}; and a case with a very light scalar in Section~\\ref{sec:dark}. Finally, we study some examples of the general case in Section~\\ref{sec:general} and conclude in Section~\\ref{sec:conclusions}. Appendix~\\ref{app:thermalpot} contains some technical details of the full one-loop analysis, including the $T=0$ renormalization conditions. ",
        "conclusions": "}  Our analysis of strong EWPhTs in the SM plus singlet highlights the richness of possibilities this simple extension of the SM offers. By relying on a simple mean-field approximation to the finite temperature scalar potential and a judicious choice of parametrization we have been able to perform a thorough analytical study of strong EWPhTs triggered by tree-level dynamics. We have given a strategy, summarized in Table~\\ref{Table1}, to identify the regions of parameter space that would lead to such strong EWPhTs. At the same time, our analytical approach has improved the understanding of the mechanisms behind such transitions allowing us to uncover new scenarios not appreciated before. One interesting example are those transitions that rely on the presence of a flat direction in the potential at the critical temperature $T_c$, mechanism that could operate in many other models besides the SM plus singlet one.  We have computed the important ratio $v(T_c)/T_c$ both in particular realizations of the SM plus singlet model previously studied in the literature and in some representative examples of the most general model. In models with $\\mathbf{Z}_2$ symmetry we have determined that a truly strong EWPhT based on a tree-level barrier must proceed from an EW symmetric vacuum that breaks the $\\mathbf{Z}_2$ symmetry to an EW broken vacuum that is $\\mathbf{Z}_2$ symmetric. In other particular models and in general, we have shown that $v(T_c)/T_c$  can be easily larger than 1. This is a necessary requirement for  successful electroweak baryogenesis (switching off in the broken phase sphaleron processes that would erase the created baryon asymmetry). This jump in $v(T)$ is necessarily associated with a jump in the singlet VEV, which can also be relevant in some baryogenesis mechanisms. In addition such strong EWPhTs could also lead to a relic stochastic background of gravitational waves.  We have refined our analysis going beyond the mean-field approximation by including one-loop effects, with appropriate renormalization conditions at $T=0$ and inclusion of thermal effects with daisy resummation and no high-temperature expansions. Although $v(T_c)/T_c$ is lowered by such refinement one still finds strong EWPhTs.  As a byproduct, our parametrization of the scalar potential of the SM plus singlet, which allows a good control over its vacuum structure, can be useful also at $T=0$. In fact, the conditions summarized in the lower part of Table~\\ref{Table1} can be applied at $T=0$ to guarantee the stability of the EW vacuum and then our parametrization can be applied to phenomenological analyses allowing a direct control over physical quantities.  \\newpage \\appendix"
    },
    "1107/1107.2920_arXiv.txt": {
        "abstract": "We present the analysis of 4 months of {\\it Kepler} photometry of the K4V star HAT-P-11, including 26 transits of its ``super-Neptune'' planet. The transit data exhibit numerous anomalies that we interpret as passages of the planet over dark starspots. These spot-crossing anomalies preferentially occur at two specific phases of the transit. These phases can be understood as the intersection points between the transit chord and the active latitudes of the host star, where starspots are most abundant. Based on the measured characteristics of spot-crossing anomalies, and previous observations of the Rossiter-McLaughlin effect, we find two solutions for the stellar obliquity $\\psi$ and active latitude $l$: either $\\psi = 106^{+15}_{-11}$ and $l=19.7^{+1.5}_{-2.2}$, or $\\psi = 97^{+8}_{-4}$ and $l=67^{+2}_{-4}$ (all in degrees). If the active latitude changes with time in analogy with the ``butterfly diagram'' of the Sun's activity cycle, future observations should reveal changes in the preferred phases of spot-crossing anomalies. ",
        "introduction": "We have been developing a new method to measure the obliquity of a star with respect to the orbital plane of a transiting planet. Obliquities are important because they are fundamental geometric parameters, and because they bear clues about the formation, migration, and tidal evolution of close-in planets (see, e.g., Queloz et al.\\ 2000; Ohta et al.\\ 2005; Winn et al.\\ 2005, 2010a; Fabrycky \\& Winn 2009; Triaud et al. 2010; Morton \\& Johnson 2011). The traditional method involves observations of the Rossiter-McLaughlin (RM) effect, a spectroscopic anomaly that occurs during transits due to selective blockage of the stellar rotation field (see, e.g., Queloz et al.\\ 2000; Ohta et al.\\ 2005; Gaudi \\& Winn 2007). The new method is purely photometric, and is based on observations of photometric anomalies in transit light curves resulting from the passage of the planet in front of starspots.  When the planet blocks the light coming from a relatively dark portion of the stellar photosphere, the fractional loss of light is temporarily reduced, and a positive ``bump'' is observed in the light curve (see, e.g., Rabus et al.~2009). Sanchis-Ojeda et al.~(2011) showed how the recurrence (or not) of these starspot-crossing anomalies can be used to measure or place bounds upon the stellar obliquity. In simplest terms, the recurrence of an anomaly in two closely-spaced transits is evidence for a low obliquity, because otherwise the spot would rotate away from the transit chord. We and our colleagues applied this technique to the particular system WASP-4, showing that the new method gives stronger constraints on the stellar obliquity than the previous observations of the RM effect (Triaud et al.~2010).  Independently, Nutzman et al.~(2011) used the recurrences of spot-crossing anomalies (as well as the phase of the out-of-transit modulation of the total light) to confirm a low obliquity for CoRoT-2b. More recently, D{\\'e}sert et al.~(2011) found a similar pattern of recurrences for Kepler-17b, and concluded that the host star has a low obliquity.  The pattern in that case was made even more dramatic by the near-integral ratio between the rotational and orbital periods.  Sanchis-Ojeda et al.~(2011) noted that another interesting target for this method would be HAT-P-11, a K4V star with a ``super-Neptune'' planet of mass 26~$M_\\oplus$ and radius 4.7~$R_\\oplus$ in a 4.9-day orbit. The star was already suspected of having starspots (Bakos et al.~2010), it had been shown to have a high obliquity based on RM observations (Winn et al.~2010b, Hirano et al.~2011), and most importantly, it lies within the field of the view of the {\\it Kepler} photometric satellite (Borucki et al.~2010). Several months of nearly continuous, highly precise {\\it Kepler} photometry are already available, and several years of data will eventually become available.  In this paper we analyze public data from {\\it Kepler}, which have indeed revealed numerous spot-crossing anomalies, and led to constraints on the stellar obliquity, although not in the manner we anticipated.  Section 2 presents the data, and Section 3 gives our estimates for the basic system parameters as well as the times of spot-crossing anomalies. The anomalies occurred preferentially at two particular phases of the transit. Section 4 presents a simple geometric model, in which special phases are the intersection points between the transit chord and the active (starspot-rich) latitudes of the star. Section 5 discusses the results.  While preparing this paper we learned that two other analyses of the public {\\it Kepler} data have been undertaken, by Southworth~(2011) and by Deming et al.~(2011). We refer the reader to those works for a different perspectives on the data, focusing on refinement of the transit parameters rather than the spin orientation of the star. ",
        "conclusions": "\\begin{deluxetable*}{lcccc} \\tabletypesize{\\scriptsize} \\tablecaption{Results of MCMC analysis\\label{tbl:mcmc}} \\tablewidth{0pt}  \\tablehead{ \\colhead{Parameter} & \\colhead{Edge-on solution} & \\colhead{Edge-on solution} & \\colhead{Pole-on solution} & \\colhead{Pole-on solution} \\\\ \\colhead{} & \\colhead{} & \\colhead{($+$~RM prior)} & \\colhead{} & \\colhead{($+$~RM prior)} }  \\startdata Projected obliquity, $\\lambda$~[deg]  & $90\\pm 28$ & $106^{+15}_{-12}$  & $83^{+77}_{-65}$&   $121^{+24}_{-21}$ \\\\ Stellar inclination, $i_s$~[deg]  & $80^{+5}_{-3}$  & $80^{+4}_{-3}$   & $160^{+9}_{-19}$  & $168^{+2}_{-5}$ \\\\ Obliquity, $\\psi$~[deg]  & $91  \\pm 27$ &  $106^{+15}_{-11}$ &  $89^{+27}_{-27}$ & $97^{+8}_{-4}$ \\\\ Latitude of active zone~[deg]  & $19.3^{+1.7}_{-3.0}$  & $19.7^{+1.5}_{-2.2}$ & $63^{+6}_{-17}$  & $67^{+2}_{-4}$  \\\\ Half-width of active zone~[deg] & $4.5^{+1.7}_{-2.1}$ &$4.8^{+1.5}_{-1.8}$  & $4.0^{+2.3}_{-2.2}$  & $4.5^{+1.6}_{-1.9}$ \\enddata \\end{deluxetable*}  The main results of our study are (1) the finding that the starspots on HAT-P-11 are preferentially found at certain active latitudes; (2) the confirmation that the HAT-P-11 star is misaligned with the orbit of its close-in planet; (3) the placement of quantitative bounds on the three-dimensional stellar obliquity based on the observed pattern of spot anomalies and a simple geometrical model. Two families of geometric solutions were found, one in which the star is viewed nearly pole-on, and the other in which the star is viewed closer to edge-on.  Even though both solutions agree on the stellar obliquity, they differ in the locations of the active latitudes: one solution gives more Sun-like active latitudes of about $20^\\circ$ while the other solution favors a larger latitude of about $60^\\circ$. Breaking this degeneracy would therefore tell us whether or not the activity cycle on HAT-P-11 resembles that of the Sun in this respect.  One might think that the measurement of $v\\sin{i_s}$ would help to distinguish these two cases, but the two solutions predict values of $v\\sin{i_s} = (2\\pi R_\\star/P_{\\rm rot}) \\sin{i_s}$ equal to $1.3$ km~s$^{-1}$ and $0.5$~km~s$^{-1}$, which are both compatible with the observed value of $1.5 \\pm 1.5$ km~s$^{-1}$ (Bakos et al.~2010). A basic difference between the pole-on and edge-on solutions is that the pole-on solution would produce smaller out-of-transit (OOT) flux variations for a spot of a given size and intensity.  Indeed, an exactly pole-on solution would not produce any OOT variations. For nearly pole-on configurations, the OOT variations arise only from small variations in limb darkening and geometrical foreshortening along the nearly-circular trajectory of the spot on the stellar disk.  In contrast, for the edge-on solution, the spot disappears from view for half of the rotation period, resulting in larger variations.  Our geometric model did not make use of the information borne by the out-of-transit (OOT) flux variations; could this be used to break the modeling degeneracy and sharpen the constraints?  We have made some efforts in this direction, but they have been inconclusive, mainly because the spot sizes and intensties are not known {\\it a priori}. Our simulations show that it is possible that the pole-on solution is correct, and that the $1.5\\%$ OOT variations are produced by large and/or dark spots.  Another feature of the pole-on model is that each spot crosses the transit chord twice per rotation period, which might lead to a detectable correlation between the observed spot anomalies. We explored this possibility, but could not reach a firm conclusion because a given spot may traverse the transit chord quickly enough that the planet does not always cross it. Nevertheless, future observations with {\\it Kepler} and a more detailed spot-by-spot model might one day be used to break the degeneracy.  Future observations may also reveal the changes in the active latitudes that are analogous to the equatorial drift of the active latitudes on the Sun. As is well known, over the 11-year solar activity cycle the active latitudes migrate toward the equator, where they disappear and are then recreated at higher latitudes, the phenomenon that underlies the butterfly diagram.  Should this also occur for HAT-P-11, we would observe a drift in the phases of the spot-crossing anomalies.  Assuming the active latitudes migrate toward the equator, the phases would behave differently for the pole-on and edge-on solutions. In the case of the pole-on solution, the phases will separate apart and move toward the extremes of the transit, while for the edge-on solution, the phases will move closer to one another.  Of course, we should not necessarily assume that the active latitudes migrate to lower latitudes, as is seen on the Sun, but the data themselves may be able to confirm this fact, making use of differential rotation. In fact, Messina \\& Guinan (2003) studied latitude migration on young solar analogues through differential rotation and found that for some stars the active latitudes seem to move poleward rather than migrating toward the equator. A combination of the measurement of the drift of the phases and the change in the period will give us a way of breaking the degeneracy and studying the rate of migration of the active latitudes. (For a review of starspots and the various techniques by which they are observed, see Strassmeier~2009.)  Independently of the activity cycle of HAT-P-11, the confirmation that the star has a high obliquity is helpful for understanding the origin of the close-in planet HAT-P-11b. There has been a recent resurgence of the theory that close-in planets are emplaced by few-body gravitational interactions followed by tidal capture, rather than by gradual inspiral due to interactions with a protoplanetary disk (see, e.g., Fabrycky \\& Tremaine 2007, Nagasawa et al.~2008, Matsumura et al.~2010). Part of the evidence is the large incidence of high obliquities among stars that host close-in giant planets (see, e.g., Triaud et al.~2010). With our corroboration of the RM results by Winn et al.~(2010b) and Hirano et al.~(2011), HAT-P-11 can be firmly added to this roster."
    },
    "1107/1107.0608_arXiv.txt": {
        "abstract": "We present a novel method for combining the analog and photon-counting measurements of lidar transient recorders into reconstructed photon returns. The method takes into account the statistical properties of the two measurement modes and estimates the most likely number of arriving photons and the most likely values of acquisition parameters describing the two measurement modes. It extends and improves the standard combining (``gluing'') methods and does not rely on any \\emph{ad hoc} definitions of the overlap region nor on any background subtraction methods. ",
        "introduction": "In order to extend the total dynamic range of measurements the same back-scattered return signal in modern lidar acquisition systems (a.k.a.\\ transient recorders) is usually sampled with two distinct methods: a fast analog-to-digital converter and a photon counting unit \\cite{eichinger}. The two range-resolved traces are then combined (``glued'') by first correcting the dead-time effects of the photon counting mode \\cite{donovan} and then by calibrating (fitting) the analog trace to the photon counts in some suitable photon-counting rate interval \\cite{gluing,newsom,whiteman}.  The final step in the construction of the ``glued'' trace involves choosing a suitable signal size above which only rescaled analog values are considered and below which only the photon-counting trace is used.  The general usability of such ``gluing'' methods is hampered by several intrinsic weaknesses: the ``background'' is usually subtracted from both measurement modes whereas it could be retained and used as additional information; a large variety of regressions and $\\chi^2$ minimizations are used in the calibration of the analog signal; arbitrary and not well defined photon-counting rate fitting intervals are imposed in order to stabilize the former minimizations; photon-counting nonlinearity is usually corrected only with manufacturer-supplied dead-time values \\cite{licel-gluing}; the unused half of the measured data is simply discarded and the actual position of the crossover between the analog and the photon-counting trace is not selected by any strict rules. ",
        "conclusions": "Modern transient recorders offer two principally different measurements of the same photon return: the digitization of the analog signal and the photon-counting mode. We are thus challenged, not only to use them in their respective regions of validity (like the usual ``gluing'' methods) but to combine them into a more accurate estimate of the photon numbers by using detailed statistical models of both acquisition processes.  The maximum likelihood procedure described in this work offers a reconstruction of photon returns that has a natural transition between the analog and photon-counting signals and is based on their analytical measurement models. In this work we have been using fairly rudimentary models of the two measurement processes, nevertheless they still adequately capture the main strengths of the new reconstruction procedure. Furthermore, if more detailed models are needed, they can be simply included into the probability expressions entering the likelihood function.  In this method we are strongly discouraged to attempt any kind of background subtraction. The background photons are treated as a normal signal since they appear in both measurement modes as viable data. Any removal of background (from dawn or daylight) should be done on the final reconstructed photon numbers.  The maximum likelihood method works in all conditions, even in the presence of clouds or other enhanced scattering objects. It fails only when the input levels of photons are not exploring the whole dynamic range of the transient recorder (i.e.\\ both corners of Fig.~\\ref{f:am}) and thus no reliable estimate of the acquisition parameters $\\alpha$, $\\beta$, $\\gamma^2$, and $\\delta$ can be obtained. Through the offset analysis of the likelihood value it offers a simple way for estimation of the potential relative delay between the two measurement traces, which in our particular case turned out not to be negligible.  The code implementing the maximum-likelihood reconstruction of lidar returns is available under the GPL3 license at \\href{http://www.ung.si/~darko/lidar/}{http:/\\!\\!/www.ung.si/\\textasciitilde darko/lidar/}."
    },
    "1107/1107.0004_arXiv.txt": {
        "abstract": "We present a weak-lensing analysis of the merging cluster A2163 using Subaru/Suprime-Cam and CFHT/Mega-Cam data and discuss the dynamics of this cluster merger, based on complementary weak-lensing, X-ray, and optical spectroscopic data sets. From two-dimensional multi-component weak-lensing analysis, we reveal that the cluster mass distribution is well described by three main components including a two component main cluster A2163-A with mass ratio 1:8, and its cluster satellite A2163-B. The bimodal mass distribution in A2163-A is similar to the galaxy density distribution, but appears as spatially segregated from the brightest X-ray emitting gas region. We discuss the possible origins of this gas-dark matter offset, and suggest the gas core of the A2163-A subcluster has been stripped away by ram pressure from its dark matter component. The survival of this gas core to the tidal forces exerted by the main cluster lets us infer a subcluster accretion with a non-zero impact parameter. Dominated by the most massive component of A2163-A, the mass distribution of A2163 is well described by a universal Navarro-Frenk-White profile as shown by a one-dimensional tangential shear analysis, while the singular-isothermal sphere profile is strongly ruled out. Comparing this cluster mass profile with profiles derived assuming intracluster medium hydrostatic equilibrium (H.E.) in two opposite regions of the cluster atmosphere has allowed us to confirm the prediction of a departure from H.E. in the eastern cluster side, presumably due to shock heating. Yielding a cluster mass estimate of $M_{500}=11.18_{-1.46}^{+1.64}\\times10^{14}h^{-1}\\Msun$, our mass profile confirms the exceptionally high mass of A2163, consistent with previous analyses relying on the cluster dynamical analysis and $Y_{\\rm X}$ mass proxy. ",
        "introduction": "Galaxy clusters are the largest self-gravitating systems in the universe, residing at the intersection of large-scale filamentary structures of the cosmic web. According to the hierarchical structure formation scenario based on a cold dark matter (CDM) paradigm, mass accretion flows onto clusters are still ongoing and high-mass ratio cluster mergers, the so-called major mergers, sometimes occur. Major cluster mergers are among the most energetic events in the universe, releasing amounts of gravitational energy as large as $10^{64-65} {\\rm erg}$. Spatially resolved X-ray observation has revealed to us that this collision energy partly dissipates in the intracluster medium (ICM) through shock heating and turbulence, yielding a complex ICM brightness and thermal structure \\citep[see, e.g.,][and reference therein]{mar07}. The dynamics of cluster collisions are however dominated by the cluster mass distribution, which cannot be constrained from X-ray observations alone. Due to the collisional nature of ICM, cluster mergers are indeed expected to generate some transient decoupling between spatial distributions of the CDM and hot X-ray emitting gas. Moreover, mass estimates relying on gas properties may be biased by merger-induced perturbations of the ICM hydrostatic equilibrium (H.E.) within individual colliding clusters.  Weak-lensing distortions of background galaxy images provide us with a unique opportunity to reconstruct the distribution of matter in clusters without any assumption of mass model and dynamical states \\citep[e.g., ][]{kai95, bar01, sch06, oka08, oka10b} and measure cluster masses \\citep[e.g., ][]{gra02, gav03, bar07, hoe07, oka10b, ume09}. Therefore, a weak-lensing analysis of merging clusters \\citep[e.g.,][]{clo06,mah07,oka08,oka10b,mer11} is an observational breakthrough in measuring mass distribution, thereby providing complementary information to X-ray measurements. \\cite{oka08} conducted a systematic study of seven merging clusters, representing various merging stages and conditions, based on a joint weak-lensing, optical photometric, and X-ray analysis, and revealed that the mass and optical light of member galaxies are similarly distributed in merging clusters regardless of their merging stages, but the mass distribution in merging clusters is highly irregular, and is quite different from the ICM distributions. It indicates that the gas and mass evolutions under cluster mergers are different. This feature is confirmed by weak-lensing studies of 30 clusters as a collaboration of ``Local Cluster Substructure Survey'' \\citep{oka10b}. Thus, a joint analysis of clusters \\citep[e.g.,][]{mah08,oka08,oka10c,kaw10,ume10,zha10} will yield a comprehensive and quantitative understanding of cluster merger physics involved in the structure formation.  Located at a redshift of $0.203$,  A2163 is a rich, X-ray luminous and hot galaxy cluster showing various signatures of ongoing merger events, including irregular optical and X-ray morphologies \\citep[see, e.g.][]{elb95,mar01,mau08,bou10}, and a prominent radio halo emission \\citep{fer01,fer04}. This cluster has been known to exhibit a clear spatial segregation between a bimodal galaxy distribution and a more centrally peaked ICM morphology \\citep[][hereafter M08]{mau08}, while the projected distance separating its main mass and gas centroids has been measured as one of the largest in a sample of 38 clusters analyzed from both strong galaxy lensing and X-ray imaging \\citep[][]{shan10}. As revealed from spectroscopic and photometric analyses in M08, the galaxy distribution in A2163 can be separated into two components: the massive cluster A2163-A and its northern companion A2163-B, and the main cluster component itself showing a bimodal morphology with two brightest galaxies (BCG1 and BCG2). In a more recent analysis of \\textit{Chandra} and \\textit{XMM-Newton} data, \\citet[][hereafter B11]{bou10} evidenced the westward motion of a cool core across the E-W elongated  atmosphere of the main cluster, A2163-A. Located close to the second galaxy over-density, this gas 'bullet' appears to have been spatially separated from its galaxy component as a result of high-velocity accretion. From gas brightness and temperature profile analysis performed in two opposite regions of the main cluster, B11 further showed that the ICM has been adiabatically compressed behind this crossing ``bullet''.  Characterized by an exceptionally high ICM temperature first measured in X-ray \\citep[$k_BT=13.9$~keV;][]{arn92} then confirmed from its Sunyaev-Zeldovich (SZ) distortion \\citep{Nord09}, A2163 has been suggested to be exceptionally massive from various X-ray and SZ analyses \\citep[see, e.g.][]{vik09,Planck11} via the mass scaling relation. Former weak-lensing analyses have been conducted using the Very Large Telescope and CFHT/Mega-Cam analyses of \\citet[][]{cyp04} and \\citet[][]{rad08}. These analyses, however, did not exclude merger galaxies in the background shear catalog, as demonstrated by other studies \\citep{bro05,oka08,ume08,ume09,oka10b,ume10}. As shown by \\cite{oka10b}, a contamination of member galaxies significantly dilutes lensing distortion signals, mainly in the central region, thereby yielding biased estimations on cluster parameters.   In this paper, we conducted a weak-lensing analysis using Subaru/Suprime-Cam and made a secure selection to avoid a contamination of member galaxies in the shear catalog, combining CFHT/Mega-Cam data. The structure of this paper is as follows:  in Section~\\ref{sec:data} we briefly describe the weak-lensing analysis. The projected distributions of mass, member galaxies and hot-gas are compared in Section \\ref{sec:map}. Section \\ref{sec:1d} is devoted to a one-dimensional shear analysis to measure cluster mass. In Sec \\ref{sec:2d}, we perform, based on a two dimensional shear pattern, a multi-components fitting to measure three components revealed by M08. We compare weak lensing mass with dynamical and X-ray hydrostatic equilibrium masses in Section ~\\ref{sec:mass}. In Section ~\\ref {sec:dis}, we discuss quantitatively the physical process of making an offset between mass and gas distributions. Section ~\\ref{sec:sum} summarizes our results. The cosmological parameters of $H_0 = 100h^{-1}$ Mpc$^{-1}$ km s$^{-1}$, $\\Omega_0=0.3$, and $\\Lambda_0=0.7$ are used in this paper. Given the cosmology, $1\\farcm=140.21~h^{-1}{\\rm kpc}$. ",
        "conclusions": "\\label{sec:dis}      From the inversion of a background galaxy shear pattern, we revealed a bimodal mass distribution in the central region of A2163 and various anisotropies at larger radii, the most significant of which corresponding to the northern subcluster A2163-B. Modeling the underlying three dimensional distribution of the cluster mass with three components further allowed us to constrain the mass ratio between the two central components to 1:8, and attribute comparable mass values to the central and northern subclusters.  The spatial coincidence between X-ray peak, mass and member galaxy distribution in A2163-B suggests  that A2163-B did not yet interact with A2163-A.  While coinciding with the bimodal galaxy distribution revealed in M08, the central mass distribution appears instead as spatially segregated from the brightest X-ray emitting gas region. In particular, the gas bullet suggested to cross the cluster atmosphere from the evidence of a cold front, in B11, appears as preceded by the secondary mass peak to a projected distance of $\\sim2\\farcm\\sim280h^{-1}~{\\rm kpc}$ along its westward direction of motion. Assuming an infalling subcluster has experienced a transient separation of its gas and dark matter components as a result of ram pressure stripping, this configuration would require the gas core to have survived tidal distortions exerted by the main cluster, and the core temperature to be consistent with the subcluster mass. In the following, we discuss how these conditions allow us to put some additional constraints on the kinematics of the ongoing subcluster accretion.  \\subsubsection*{The Ram Pressure Stripping Condition}  One possibility for explaining the offset between mass and gas distributions is the ram-pressure stripping: if the ram-pressure force on the X-ray core is stronger than the gravity around the core region, the gas is stripped from its gravitational potential. This condition is expressed as \\begin{eqnarray} \\frac{G M(< r_{\\rm core}) \\rho_{\\rm core}}{r_{\\rm core}^2} <  A (\\pi r_{\\rm core}^2 \\rho_{\\rm sur} v^2) \\left(\\frac{4}{3} \\pi r_{\\rm core}^3\\right)^{-1} \\label{eq:Pram} \\end{eqnarray} \\citep{tak06}. Here, $r_{\\rm core}$ is the radius of the X-ray core, $M(< r_{\\rm core})$ is the spherical mass of the subcluster within the radius $r_{\\rm core}$, $\\rho_{\\rm core,sur}$ are the density of the core and its surrounding gas, respectively, and $v$ is the velocity of the core in the center-of-mass frame. $A$ is a fudge factor at an order of unity. This is why the ram-pressure stripping is not effective due to a Kelvin-Helmholtz instability, magnetic fields, and shock heating. B11 shows a lower limit of the density jump, $\\rho_{\\rm core}/\\rho_{\\rm sur}=1.28$, at the west edge of the core (red-curved line in Figure \\ref{fig:chan+kappa}).     We consider the ram-pressure stripping condition on the sub cluster. The substructure mass inside $r_{\\rm core}$ is calculated from the TSIS model obtained by the weak-lensing analysis.  Since $M_{\\rm TSIS} \\propto r$, the condition (\\ref{eq:Pram}) is free from the core size. Although we consider the internal structure of gas, the resultant condition does not significantly change. Figure \\ref{fig:Pram} plots the ram-pressure stripping condition $P_{\\rm ram}/P_{\\rm grav}-1$ as a function of infall velocity $v$, where $P_{\\rm ram}/P_{\\rm grav}$ gives the ratio of the right-hand side to the left-hand side of Equation (\\ref{eq:Pram}). We adopt $A=0.5~{\\rm and}~0.9$. M08 found the gradient of line-of-sight velocity $v_{\\rm los}\\sim1250~{\\rm km~s^{-1}}$, giving the lower-limit of infall velocity. The infall velocity is expected from the masses to have an order of \\begin{eqnarray} v\\sim \\left[\\frac{2G(M_{\\rm main}+M_{\\rm sub})}{r_{\\rm main}+r_{\\rm sub}}\\right]^{1/2}\\sim 2100 ~{\\rm km~s^{-1}} \\label{eq:vmerger} \\end{eqnarray} \\citep{ric01}, where $M$ and $r$ are the virial mass and radius, respectively. We here used observed truncated mass and radius for the sub cluster. The case of $A=0.9$ satisfies the ram-pressure stripping condition for all velocities above $\\delta v_{\\rm los}$. The case of $A=0.5$ also satisfies the condition around $2100~{\\rm km~s^{-1}}$. Therefore, the  X-ray core initially associated with the sub cluster is easily stripped away from its central region. However, we must keep in mind that the estimated density jump might give the lower limit, due to the deprojection method we applied. In this case, a requirement of the ram pressure condition gives $\\rho_{\\rm core}/\\rho_{\\rm sur}<5(A/1)(v/2100~{\\rm km~s^{-1}})^2$.      \\subsubsection*{Constraints on impact parameter from survival condition}  We next discuss the destruction process of the gas core by the strong tidal field of the main cluster. Numerical simulation \\citep{tor04} has shown that the lifetime for a gas satellite is shorter than that for a dark matter one. In particular, the lifetime for a gas core decoupled from dark matter potential is much shorter than that of a dark matter satellite.  We roughly estimate the gas core disruption by the tidal force of the main cluster. As shown above, the gas core could easily be stripped away from the subcluster's core region. We therefore consider % the stripped core which is composed of the gas, to be free from the gravity of the subcluster. The gas model for the core is adopted as a single power-law density profile as an approximation, \\begin{eqnarray} n_e(r)=n_{e,0} \\left(\\frac{r}{r_0}\\right)^{-\\alpha}, \\end{eqnarray} where the normalization $n_{e,0}$ and the slope $\\alpha$ are the model parameters. We use the position of cold front $r_0=0\\farcm6$ (red-curved line of the right panel of Figure \\ref{fig:chan+kappa}). The tidal radius at which the core density is truncated is estimated by the force equivalence between the internal gravity and external tides \\citep{tor98}, \\begin{eqnarray} r_t=\\tilde{A} b\\left[\\frac{M_{\\rm core}(<r_t)}{(2-\\partial \\ln M_{\\rm main}/\\partial \\ln b) M_{\\rm main}(b)}\\right]^{1/3}, \\label{eq:tidal} \\end{eqnarray} in the limit of $r_t\\ll b$ and $M_{\\rm core}\\ll M_{\\rm main}$. Here $r_t$ is the tidal radius and $b$ is the core position from the center of the main cluster; that is, an impact parameter. $\\tilde{A}$ is a fudge factor because hydrodynamic instabilities effectively destruct the gas core. We assume $\\tilde{A}=1$ in order to consider the tidal disruption only. We also assume that the density profile within a tidal radius dose not modify before and after the tidal stripping. The main cluster mass profile uses the best-fit NFW model in the case of NFW+TSIS+NFW (Table \\ref{tab:mass2}). We first calculate the minimum tidal radius $r_{t,{\\rm min}}$ as a function of the density normalization and its slope and impact parameter. The left panel of Figure \\ref{fig:taidal} shows a distribution of the minimum tidal radius of the gas core in the parameter plane of the density normalization and the slope. The contours denote the minimum tidal radius in the range of $0\\farcm02-0\\farcm1$. The point denotes the normalization and slope parameters for the core XW at the cold front. In the parameter plane, the minimum tidal radius is much smaller than the observed radius $r\\sim0\\farcm6\\sim120h_{70}^{-1}~{\\rm kpc}$, indicating that we cannot observe a remnant of the gas core if A2163-A is a system of close encounter.   We next try to constrain the impact parameter, $b$, so as to realize the observed radial size of the core XW ($r_t=r_0$) through Equation. (\\ref{eq:tidal}). The right panel shows a distribution of the impact parameter in the parameter plane of the density normalization and the slope. The contours present the impact parameters $(0.2-0.4)r_{\\rm vir}$ to explain the core size. The observed X-ray core requires the impact parameter $b\\sim0.25 r_{\\rm vir}\\sim560~{\\rm kpc}h^{-1}$. If a subcluster collides into the main cluster with this impact parameter, the central region of the main cluster is significantly affected by cluster merger. It does not conflict with the disturbed core structures and the complex temperature distribution (Figure \\ref{fig:chan+kappa}). Since this estimation is based on the gravitational process only, we need to keep in mind a possibility that hydrodynamic instabilities shorten the lifetime of gas. We emphasize that this method demonstrates only one of the methods for joint X-ray and weak-lensing analysis.            \\subsubsection*{Temperature Comparison}    Assuming the cool gas core XW and the secondary mass peak MW represent the separated components of a formerly accreted subcluster, the intrinsic temperature of XW should agree with the expectation of a self-similar subcluster temperature corresponding to the mass of MW. Given our multi-component mass distribution (NFW+TSIS+NFW; see Table 2), we discuss the consistency of this hypothesis with projected ICM temperatures derived near the cool gas core XW in B11.  The cool gas core XW being located a relatively close projected distance from the main cluster emission peak XC, average ICM temperatures measured along a line of sight intercepting XW would consist of a linear combination of the main and subcluster temperatures. These temperatures can be predicted within the radial range [0.2-0.5]$~r_{\\mathrm{500}}$, from the $M_{500}-T_{[0.2-0.5]~r_{500}}$ scaling relation calibrated by \\cite{oka10c}, in a sample of 12 local clusters analyzed in both X-ray and weak-lensing. To perform this estimation, we first measure the main cluster mass, $M_{500,{\\rm main}}=7.59_{-1.57}^{+1.22}\\times10^{14}h^{-1}\\Msun$, by multi-component analysis (NFW+TSIS+NFW), within the overdensity radius $r_{500}$. We further determine the subcluster mass, $M_{500,{\\rm sub}}$, from its current mass, $M_{\\rm sub} = M_{\\rm main}/8$, assuming the subcluster mass profile to have followed a universal NFW radial distribution before the ongoing accretion, and the profile concentration to be related to $M_{sub}$ by the halo mass-concentration dependence of \\cite{duf08}. These assumptions would yield a cluster and subcluster temperature prior to the accretion of $k_BT = 10.1^{+1.0+1.1+1.8}_{-1.4-1.0-1.8}~\\mathrm{keV}$ and $k_BT = 3.7^{+0.5+0.3+1.2}_{-0.9-0.3-1.2}~\\mathrm{keV}$, respectively, where the first, second and third errors are from mass measurement error, normalization error of scaling relation, and intrinsic scatter in the relation. Assuming the main and subcluster emissivities to be comparable from a rough analysis of the X-ray image and weighting these temperatures following a scheme proposed in \\citet{maz04}, an average ICM temperature measured in the direction of the cool core would reach a value of $k_BT \\sim 5.7~\\mathrm{keV}$. Despite being consistent with uncertainties in our mass estimates, this value is considerably lower than projected temperatures measured in % the cool gas core XW in the \\textit{XMM-Newton} temperature map and the \\textit{Chandra} temperature profile of B11 ($k_BT \\simeq 9\\pm1~\\mathrm{keV}$). Among the possible origin for this discrepancy, the cool gas core might have been heated from its virial temperature while crossing the main cluster atmosphere, possibly due to mixing with the main cluster atmosphere or a reverse shock, presumably at an earlier stage of its accretion and prior to the formation of the cold front.  \\subsubsection*{}   In summary, the density of the main cluster atmosphere and the velocity of its presumably free-falling subcluster are large enough to have separated the subcluster gas and dark matter components through ram-pressure stripping. Moreover, the survival of the subcluster gas core against tidal forces exerted by the main cluster implies the subcluster accretion to have occurred with a non-zero impact parameter. Comparing the average X-ray temperature in the direction of the cool core to its expectation from projection of the cluster and subcluster virial temperatures yields however a mild inconsistency, suggesting the cool core was partially heated while crossing the main cluster atmosphere. Deeper X-ray observations should help better constrain the shape, temperature and mass of the cool core and refine this accretion scenario."
    },
    "1107/1107.0142_arXiv.txt": {
        "abstract": "We investigate the Faraday rotation measure (RM) due to the intergalactic magnetic field (IGMF) through the cosmic web up to cosmological distances, using a model IGMF based on turbulence dynamo in the large-scale structure of the universe. By stacking the IGMF and gas density data up to redshift $z=5$ and taking account of the redshift distribution of polarized background radio sources against which the RM is measured, we simulate the sky map of the RM. The contribution from galaxy clusters is subtracted from the map, based on several different criteria of X-ray brightness and temperature. Our findings are as follows. The distribution of RM for radio sources of different redshifts shows that the root-mean-square (rms) value increases with redshift and saturates for $z \\ga 1$. The saturated value is RM$_{\\rm rms} \\approx$ several ${\\rm rad~m^{-2}}$. The probability distribution function of $|{\\rm RM}|$ follows the lognormal distribution. The power spectrum has a broad plateau over the angular scale of $\\sim 1 - 0.1^\\circ$ with a peak around $\\sim 0.15^\\circ$. The second-order structure function has a flat profile in the angular separation of $\\ga 0.2^\\circ$. Our results could provide useful insights for surveys to explore the IGMF with the Square Kilometer Array (SKA) and upcoming SKA pathfinders. ",
        "introduction": "\\label{s1}  The cosmic web of filaments and clusters of galaxies, which is predicted in the highly successful $\\Lambda$CDM cosmology \\citep{bkp96}, is filled with ionized plasma, {\\it i.e.}, the intergalactic medium (IGM) \\citep{co99,krcs05}. The hot gas with $T > 10^7$ K is found mostly in the intracluster medium (ICM) and cluster outskirts, and the gas with $10^5$ K $ < T < 10^7$ K, which is also referred as the Warm Hot Intergalactic Medium (WHIM), is distributed mostly in filaments. They were heated mostly by cosmological shock waves which were formed in the course of the large-scale structure (LSS) formation of the universe \\citep{rkhj03,psej06,krco07,skillman08,hoeft08,vazza09}. The diffuse gas with $T < 10^5$ K resides mainly in sheetlike structures and voids.  The IGM is expected to be permeated by magnetic fields just as the interstellar medium within galaxies. Theoretical studies have predicted the existence of the intergalactic magnetic field (IGMF). A number of mechanisms that can create seed magnetic fields have been suggested; besides processes based on inflation and phase transitions \\citep[see, e.g.,][for review]{wrss10} and plasma physical processes \\citep[see, e.g.,][for review]{rstt10} in the early universe, astrophysical processes include the field generations during the reionization of the universe \\citep{gfz00,lap05}, by first stars \\citep{xu08,ads10}, and at cosmological shock waves \\citep{kcor97,rkb98}, as well as the leakage of magnetic fields and cosmic ray particles from galaxies \\citep{ddlm09,mb11}. The seed fields can be further amplified by flow motions induced by the hierarchical clustering during the LSS formation \\citep{kcor97,rkcd08,sbsa10}.  Faraday rotation, the rotation of the plane of linearly-polarized radio emission due to the birefringence of magneto-ionic media, provides an observational means to exploring the IGMF. There have been a number of studies of the intracluster magnetic field (ICMF) through observations of rotation measure (RM) \\citep[see][for a review]{ct02}. For instance, RMs of hundreds ${\\rm rad\\ m^{-2}}$ were observed in clusters, indicating an average strength of the ICMF to be $\\sim 1$--10\\ $\\mu$G \\citep{ckb01,cla04,gov10}. In addition, RM maps of clusters were analyzed to get the power spectrum of turbulent magnetic fields in the ICM; for instance, a Kolmogorov-like spectrum with a bending at a few kpc scale was found in the cooled core region of the Hydra cluster \\citep{ve05}, and spectra consistent with the Kolmogorov spectrum were reported in the wider ICM for the Abell 2382 cluster \\citep{gmgp08} and for the Coma cluster \\citep{bfmg10}.  Studies of RM outside clusters, through the cosmic web, are, on the contrary, still scarce \\citep[e.g.,][]{xkhd06}, because detecting small RM is difficult with current observational facilities, and also removing the galactic foreground is not a trivial task. Nevertheless, recently, constraints for the RM through the LSS have been discussed by several authors \\citep{tss09,mao10,sch10,sts11}. For instance, \\citet{sch10} argued that in the catalog of \\citet{tss09}, the extragalactic contribution to the width of RM distribution would be $\\sigma_{\\rm RM,EG}\\sim 6$~${\\rm rad~m^{-2}}$. But the measurement error is still large, $\\sigma_{\\rm errRM}\\sim$10~${\\rm rad~m^{-2}}$; that is, so far, the nature and origin of the IGMF outside clusters has not been well constrained with RM studies. However, the next generation radio interferometers including the Square Kilometer Array (SKA), and upcoming SKA pathfinders, the Australian SKA Pathfinder (ASKAP) and the South African Karoo Array Telescope (MeerKAT), as well the Low Frequency Array (LOFAR) will enable us to investigate the IGMF outside clusters with high-sensitivity RM observations \\citep[see, e.g.,][]{cr04,beck09,kra09,glt10}.  Theoretical predictions for the RM due to the IGMF have been made with model IGMFs and simulations of the LSS formation \\citep{rkb98,dub08,cr09,ds09,sndbd10,ar10}. Especially, based on a physically motivated model, in which a part of the gravitational energy released during the LSS formation is transferred to the magnetic field energy as a result of the turbulent amplification of weak seed fields, \\citet{rkcd08} (R08) proposed that the IGMF follows largely the matter distribution in the cosmic web and the mean strength would be $\\langle B\\rangle \\sim 10$ nG in filaments in the local universe at redshift $z=0$. Studying various characteristic length scales of magnetic fields in MHD turbulence simulations, \\citet{cr09} (CR09) proposed that the RM coherence length of the IGMF in filaments would be a few $\\times\\ 100\\ h^{-1}$ kpc. Based on the model IGMF of R08, CR09 predicted that the-root-mean-square (rms) value of RM through a single filament would be of order $\\sim 1\\ {\\rm rad\\ m^{-2}}$. And using the same model IGMF, we simulated RM in the local universe \\citep{ar10} (AR10). We found that with the path length larger than the coherence length of the IGMF, the inducement of RM is a random walk process, but the resultant RM is dominantly contributed by the density peaks along the line of sight (LOS). The rms value of RM through a single filament was estimated to be $\\sim 1\\ {\\rm rad\\ m^{-2}}$, which is in agreement with the prediction by CR09. We also found that the probability distribution function (PDF) of $|{\\rm RM}|$ follows the log-normal distribution, and the power spectrum of the RM peaks at a scale of order $\\sim 1\\ h^{-1}$ Mpc.  In this paper, we extend the RM study of AR10 in the present-day, local universe by including the cosmological contribution. By taking account of the redshift evolution of the LSS and the IGMF as well as the redshift distribution of radio sources against which RM is measured, we calculate the RM through the cosmic web up to $z=5$. We then examine the spatial and redshift distributions of the RM and their statistics. In Sections 2 and 3, we describe our models results, respectively. Discussion is in Section 4, and summary and conclusion follow in Section 5. ",
        "conclusions": "\\label{s4}  Measuring Faraday RM is one of a few available methods to investigate the IGMF in the LSS of the universe, especially in filaments, which still remains largely unknown. Using a model IGMF derived from cosmological structure formation simulations, we examined the RM through the cosmic web outside of clusters, by stacking simulation data up to redshift $z=5$ and taking account of the redshift distribution of radio sources. We found that the magnitude of the RM increases with the path length; the path length and so the value of ${\\rm RM}_{\\rm rms}$ increases with the redshift depth of radio sources for $z \\la 1$ and saturates for $z \\ga 1$. The saturated value is predicted to be ${\\rm RM}_{\\rm rms} \\sim 7 - 8 \\ {\\rm rad~m^{-2}}$. We also found that the PDF of $|{\\rm RM}|$ follows the lognormal distribution. The power spectrum of the RM shows a broad plateau over the angular scale of $\\sim 1 - 0.1^\\circ$ and peaks at $\\sim 0.15^\\circ$. The second-order SF shows a flat profile in the angular separation of $\\ga 0.2^\\circ$.  Our results could provide useful insights for planning RM surveys with the SKA, and upcoming SKA pathfinders, ASKAP, MEERKAT, and LOFAR. Also we point that our results could be tested with such surveys. For instance, our power spectrum of RM peaks at the scale of $\\sim 0.15^\\circ$. A very fine RM map with the average separation of sources $\\sim 1$ arcmin, which is expected to be achieved with the SKA, should have enough angular resolution to see the power spectrum peak we predict."
    },
    "1107/1107.4282_arXiv.txt": {
        "abstract": "{ The CLEAN algorithm, widely used in radio interferometry for the deconvolution of radio images, performs well only if the raw radio image (dirty image) is, to good approximation, a simple convolution between the instrumental point-spread function (dirty beam) and the true distribution of emission across the sky. An important case in which this approximation breaks down is during frequency synthesis if the observing bandwidth is wide enough for variations in the spectrum of the sky to become significant. The convolution assumption also breaks down, in any situation but snapshot observations, if sources in the field vary significantly in flux density over the duration of the observation. Such time-variation can even be instrumental in nature, for example due to jitter or rotation of the primary beam pattern on the sky during an observation. An algorithm already exists for dealing with the spectral variation encountered in wide-band frequency synthesis interferometry. This algorithm is an extension of CLEAN in which, at each iteration, a set of $N$ `dirty beams' are fitted and subtracted in parallel, instead of just a single dirty beam as in standard CLEAN. In the wide-band algorithm the beams are obtained by expanding a nominal source spectrum in a Taylor series, each term of the series generating one of the beams. In the present paper this algorithm is extended to images which contain sources which vary over both frequency and time. Different expansion schemes (or bases) on the time and frequency axes are compared, and issues such as Gibbs ringing and non-orthogonality are discussed. It is shown that practical considerations make it often desirable to orthogonalize the set of beams before commencing the cleaning. This is easily accomplished via a Gram-Schmidt technique. ",
        "introduction": "It has been known for some decades that some radio sources vary in intensity over time (Dent \\cite{dent}). The range of time scales detected so far, roughly from minutes to decades, is set by the practicalities of measurement rather than by any intrinsic properties of the source populations. The present paper is concerned with sources which exhibit significant variation in flux density on time scales of less than 24 hours (so-called Intra-Day Variables or IDVs) and which are also associated with extended structure of a size resolvable by present-day radio interferometers. The structure makes it of interest to image such sources, but the rapid variability can hinder attempts to do so. It is these difficulties in the interferometric imaging of IDV sources which the present paper is designed to address.  How many such sources are there, and of what types? Any estimation of the incidence of variability among the radio source population is complicated by the extra degrees of freedom implicit in a non-flat light curve. It is also inevitable that there will be selection effects due to the inability of the instrument to detect modulation depths below a certain cutoff, or the insensitivity of the observing regime to time scales outside a certain range. The difference between the population of objects observable in our galaxy and those which are extra-galactic introduces a dependence of incidence on galactic latitude (further complicated by the greater propensity of compact, extragalactic sources to scintillate at low galactic latitudes). It is even more difficult to estimate the fraction of variable sources which are IDV, since surveys of the depth and cadence necessary to determine this quantity demand extravagant amounts of observing time.  Of presently-available survey results, De Vries et al (\\cite{devries}), in a two-epoch comparison of a high latitude field, found only $\\sim 0.2$ deg$^{-2}$ at 1.4 GHz flux densities of $\\ge 2$ mJy. Bower et al (\\cite{bower}) have processed a large amount of VLA 5 and 8.4 GHz archival data and find a `snapshot' incidence of $1.5$ deg$^{-2}$ radio transients at flux densities greater than 350 $\\mu$Jy. Becker et al (\\cite{becker}) compared 3 epochs of VLA observations of the galactic plane at 6 cm and found $\\sim 1.6$ deg$^{-2}$ sources between 1 and 100 mJy which had a modulation depth greater than 50\\%. Regarding the IDV fraction, Kedziora-Chudczer et al (\\cite{kedziora}) found, in their multi-frequency study (between 1.4 and 8.6 GHz), that about half of their 13 BL Lac objects, and a smaller fraction of quasars, exhibited IDV up to about 10\\% modulation; more recently Lovell et al (\\cite{lovell}), in a more comprehensive survey at 5 GHz, report that 37\\% of their 443 flat-spectrum sources showed significant variability on a 2-day time scale.  Sources which exhibit a combination of IDV plus some resolved structure fall into several classes, which are briefly reviewed in the following paragraphs.  Novae are expected to exhibit significant changes in radio flux density at intraday time scales. In the standard model, the flux density from an expanding isothermal shell is expected initially to increase proportional to the square of time since the outburst (Hjellming \\cite{hjellming}), although some recent VLBA observations appear to be inconsistent with that picture (Krauss et al \\cite{krauss}). However, most novae at least in the past decade have not been resolved by radio interferometry until of order 100 days after the outburst, by which time the fractional change in flux density has fallen to about a percent or two per day (see e.g. Eyres et al \\cite{eyres_2000}, Eyres et al \\cite{eyres_2005}, Heywood \\& O'Brien \\cite{heywood}). A recent exception is the 2006 outburst of the recurrent nova RS Ophiuchi (O'Brien et al \\cite{obrien_2006a}). This was observed with the VLBA first on day 14 after the outburst, and the European VLBI Network (EVN) from day 20 (O'Brien et al \\cite{obrien_2006b}). The size of the source at 6 cm wavelength grew from approximately 20 to 35 mas between these observations. The flux density was no longer varying rapidly by that time, but the earliest observations with MERLIN on day 4 after the outburst show it then varying by a factor of 2 over less than a day. Model calculations indicate that RS Oph would have been resolvable by the EVN at that time.  X-ray binaries form an interesting class of galactic IDV sources. Several of these are associated with resolved structure and have emitted jansky-strength radio flares which show pronounced evolution in intensity on intraday time scales. Examples include Circinus X-1 (Haynes et al \\cite{haynes}, Tudose et al \\cite{tudose_2008}), Cygnus X-3 (Tudose et al \\cite{tudose_2009}) and GRS 1915+105 (Fender et al \\cite{fender}, Rushton et al \\cite{rushton}).  Extragalactic IDV sources fall into fewer categories and are in fact dominated by flat-spectrum AGNs, or objects nowadays collectively described as blazars (Wagner \\& Witzel \\cite{wagner}, Lovell et al \\cite{lovell}). One could almost say that being a blazar is both a necessary and a sufficient condition for an extragalactic source to exhibit IDV, since these objects have the combination of brightness and small size which allows them on the one hand to be detectable at great distances and on the other to be compact enough to exhibit IDV. In many of these objects the compact core appears to be embedded in radio structure which is resolvable by current interferometers (see for example Kovalev et al \\cite{kovalev}, Ojha et al \\cite{ojha}).  IDV among blazars is an area of current interest. IDV mechanisms for these sources fall into two categories: variation which is intrinsic to the object, and variation caused by the propagation medium (scintillation in the local interstellar medium, or ISS). Both appear to play some role but the relative proportions are not clear. Spectral dependence of the variability at cm wavelengths, annual modulation of the time scales and time shifts between measurements made at different locations all offer strong evidence for the dominance of ISS for some sources (Bignall et al \\cite{bignall} and references within). Correlation between variation at radio and other wavelengths (Quirrenbach et al \\cite{quirrenbach}) or a bad fit to ISS spectral models (Fuhrmann et al \\cite{fuhrmann}) tend to support the intrinsic origin of IDV in others. Intrinsic origin presents theoretical challenges because it is difficult to explain it without invoking either an incidence of extreme Doppler factors which is hard to accept on statistical grounds, or brightness temperatures in excess of the inverse-Compton limit of about $10^{12}$ K (Jauncey et al \\cite{jauncey}). Some questions remain also in the ISS picture, such as a relatively weak correlation with galactic latitude (Kedziora-Chudczer et al \\cite{kedziora}, Lovell et al \\cite{lovell}); the connection between scintillation and scatter-broadening is also still unclear (Ojha et al \\cite{ojha}, Lazio et al \\cite{lazio}).  The only non-blazar category of extragalactic IDV source which is relevant to the present paper is the maser. McCallum et al (\\cite{mccallum}) for example show that the $\\mathrm{H}_2 \\mathrm{O}$ megamasers in Circinus are both resolvable and exhibit IDV. The techniques described in the present paper can make it easier to image IDV masers which have superimposed lines, meaning they must both be present in the same channel map.  It can thus be seen that there are many objects which it is desirable to image at radio wavelengths with as high resolution as possible, but whose variation on intraday timescales interferes with imaging via earth-rotation aperture synthesis. Broadly speaking, the problem occurs because the IDV causes the pattern of lobes around the time-variable component of the source to be poorly matched to the dirty beam. Such structure cannot be completely removed by the standard CLEAN algorithm, and will therefore tend to limit the dynamic range of the image. As further described in section \\ref{clean}, the situation is analogous to the problem in multi-frequency synthesis (MFS), caused by differences in spectral index among objects in the field. The argument for seeking an improved way to image IDV sources is the same as for this MFS case. In fact, as is shown in the present paper, the same technique, originally developed for MFS by Sault and Wieringa (\\cite{sault_wieringa}), can be applied to both situations.  The plan of the paper is as follows. Section \\ref{interferometry} contains a brief outline of aperture synthesis theory. In sections \\ref{conway} and \\ref{swa}, the generalized $N$-beam Sault-Wieringa theory is described in detail. The desirability of orthogonalizing the beams is discussed in section \\ref{orthogonality}. Different time and frequency basis functions are compared in section \\ref{basis_functions}, with emphasis on the avoidance of Gibbs-like phenomena. In section \\ref{dual}, a dual expansion in both frequency and time axes is shown to be both necessary and effective in cleaning a simulated wide-band data set which includes sources which vary significantly both in frequency and time. Finally, in section \\ref{simpler_lc_method}, a simple 2-beam expansion is derived which is shown to be a powerful means of increasing dynamic range in many IDV cases.  Brief descriptions of the methods described here have already been presented elsewhere in earlier stages of development (Stewart \\cite{stewart_2008} and \\cite{stewart_2009}). ",
        "conclusions": "In this paper, a technique developed originally by Conway et al (\\cite{conway}) and Sault and Wieringa (\\cite{sault_wieringa}) to allow cleaning of multi-frequency-synthesis images has been generalized and shown to be applicable to earth-rotation-synthesis observations in which some of the sources vary significantly in brightness over the course of the observation. Sources which vary over the course of a day are not very common, but they do occur, and are sometimes (in the case for example of novae) sources which it is of the highest interest to map accurately. But in fact one does not have to look for natural variations in flux density to encounter this problem: any movement of the primary beam of the array on the sky during an observation will generate artificial fluctuations, not only in the average flux density of sources, but, due to the frequency-dependent size of the primary beam, also in their spectral indices. Since some kind of pointing error in a mechanically tracking dish is probably unavoidable, this is likely to be a problematic issue when performing any kind of wide-field or mosaiced observation (see e.g. Bhatnagar et al \\cite{bhatnagar}), particularly so in view of current hopes for improvements in dynamic range from the several wide-band, dish-antenna arrays, such as eVLA, eMERLIN and MeerKAT, which are currently under construction.  Deconvolution of time-varying sources reveals some issues which are not usually encountered in the frequency-synthesis case. The principal one of these is that choice of basis function is now of some importance. This issue was explored with some care, in particular the avoidance of Gibbs ringing when periodic basis functions are employed. It was also shown that, provided the time variation in the field is limited to a single, unresolved source, a particularly simple 2-beam technique can produce an almost perfect deconvolution.  Whereas the original parallel decomposition treatment employed at most 3 or 4 beams, some of the situations described in the present paper required as many as 30. Use of such large numbers of beams gives rise to a computational difficulty due to the requirement in the original Sault-Wieringa algorithm to store, for $N$ input beams, of order $N^2$ cross-correlation images. It is shown here that this memory load can be much reduced if the set of beams is made orthogonal. An algorithm for doing this was proposed and it was shown how `de-orthogonalized' clean components can be recovered at the end of the cleaning process via a simple matrix inversion."
    },
    "1107/1107.3141_arXiv.txt": {
        "abstract": "{ The presence of unresolved binaries on sub-arsecond scales could explain the existence of optically thin inner holes or gaps in circumstellar disks, which are commonly referred to as ``transitional\" or ``cold\" disks, and it is the first scenario to check before making any other assumptions. } { We aim at detecting the presence of companions inside the inner hole/gap region of a sample of five well known transitional disks using spatially-resolved imaging in the near-IR with the VLT/NACO/S13 camera, which probes projected distances from the primary of typically 0.1 to 7 arcsec. The sample includes the stars  DoAr\\,21, HD\\,135344B (SAO\\,206462), HR\\,4796A, T\\,Cha, and TW\\,Hya, spanning ages of less than 1 to 10 Myr, spectral types of A0 to K7, and hole/gap outer radii of 4 to 100~AU. } { In order to enhance the contrast and to avoid saturation at the core of the point-spread function (PSF), we use narrow-band filters at 1.75 and 2.12~$\\mu$m. The ``locally optimized combination of images\" (LOCI) algorithm is applied for an optimal speckle noise removal and PSF subtraction, providing an increase of~\\mbox{0.5-1.5\\,mag} in contrast over the classic method. } {With the proviso that we could have missed companions owing to unfavorable projections,  the VLT/NACO observations rule out the presence of unresolved companions down to an inner radius of about 0\\farcs1 from the primary in all five  transitional disks and with a detection limit of  2 to 5 mag in contrast. In the disk outer regions the detection limits typically reach 8 to 9 mag in contrast and 4.7 mag for T\\,Cha. Hence, the NACO images resolve part of the inner hole/gap region of all disks with the exception of TW\\,Hya, for which the inner hole is only 4~AU. The 5$\\sigma$ sensitivity profiles, together with a selected evolutionary model, allow to discard \\emph{stellar} companions within the inner hole/gap region of T\\,Cha, and down to the \\emph{substellar} regime for HD\\,135344B and HR\\,4796A. DoAr\\,21 is the only object from the sample of five disks for which the NACO images are sensitive enough for a detection of objects less massive than $\\sim$ 13~M$_{\\rm {Jup}}$ that is, potential giant planets or low-mass brown dwarfs at radii larger than $\\sim$ 76~AU (0\\farcs63). } {These new VLT/NACO observations further constrain the origin of the inner opacity cavities to be owing to closer or lower-mass companions or other mechanisms such as giant planet formation, efficient grain growth, and photoevapo\\-ration (for DoAr\\,21 and HR\\,4796A). } ",
        "introduction": "\\begin{table*}[!thb]  \\caption{Observing log of the 2.12 $\\mu$m VLT/NACO images of the five stars with transitional disks and corresponding PSF-calibrators} \\label{table:1} \\bigskip \\scriptsize \\centering \\begin{tabular}{lcccccccccccl} \\hline\\hline \\noalign{\\smallskip} Star & Date & Grade$^{\\boldmath  \\mathrm{1}}$ & DIT & NDIT & Exp.time$^{\\boldmath  \\mathrm{2}}$ & Seeing$^{\\boldmath  \\mathrm{3}}$ & Airmass & $\\tau_{0}^{\\boldmath \\mathrm{4}}$ & Strehl$^{\\boldmath  \\mathrm{5}}$ & FWHM$^{\\boldmath  \\mathrm{6}}$ &  V$^{\\boldmath  \\mathrm{7}}$ & Notes$^{\\boldmath  \\mathrm{8}}$ \\\\ \\noalign{\\smallskip} \\& PSF &    &     & (s) &  & (min.) & (arcsec)  &  (at start) & (ms) & (\\%) & (mas) &  (mag)  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} DoAr\\,21 & 2007-06-18 &  B  &10.0 & 4 &  8.0 & 0.68   & 1.530  & 2.55 & 9.3 & 129.1  & 14.01$\\pm$0.03 & sat.  \\\\ CD-24\\,13079 &  &  & 4.0 & 1, 3 & 0.73 & 0.75  &  1.547 &  2.14  &  15.2 &  85.0  & 9.86 & not sat. \\\\ \\hline \\noalign{\\smallskip} HD\\,135344B & 2007-05-09 & B & 1.0 & 12 & 2.4 & 1.75  & 1.062  & 2.73  & 13.5 & 82.5 & 8.708$\\pm$0.017&    \\\\ HD\\,136961  & & & 1.0 & 10 & 1.67 & 1.85 & 1.071 & 2.14 & 14.5 & 84.6 &   6.754 &  \\\\ \\hline \\noalign{\\smallskip} HD\\,135344B & 2007-06-20  & B & 2.0 & 6 & 2.4 & 0.77 & 1.051 & 1.07 & 25.1 & 74.4 & 8.708$\\pm$0.017 & \\\\ HD\\,144156  & & & 1.0 & 10 &  1.67 & 0.79 & 1.076 & 1.23 & 26.1 & 72.3 & 8.50 & close comp. \\\\ \\hline \\noalign{\\smallskip} HR\\,4796A & 2007-06-14 & B & 1.0 & 10 & 2.0 &  NA  & 1.069 & 1.67 & 35.9 & 69.2 & 5.78 & \\\\ HD\\,109536  & & & 1.0 & 1 & 0.17 &  NA & 1.092 & 1.88 & 27.0 & 71.9 & 5.127 & \\\\ \\hline \\noalign{\\smallskip} T\\,Cha & 2007-05-23  & B & 25.0  & 2 & 10.0 & 0.85  & 1.735 & 1.37 & 4.8 & 116.2 & 11.86 & ellip. \\\\ CD-78\\,512  &   &  & 12.0, 2.0 & 1, 6  & 1.0, 0.8  & 1.16, 1.89  & 1.739, 1.740 & 1.28,  0.76 & 7.8, 7.6  & 97.5, 97.5 &   10.14 & ellip. \\\\ \\hline \\noalign{\\smallskip} TW\\,Hya & 2007-05-12 & A & 5.0  & 5 & 5.0 & 1.46 &  1.033 & 1.27 & 19.0 & 78.4 & 11.07 & \\\\ HD\\,94889  &  & & 1.0 & 5 & 0.83 & 1.10 & 1.053 & 0.59 & 23.9 & 75.7 & 8.80 & close comp. \\\\ \\hline \\noalign{\\smallskip} TW\\,Hya & 2007-05-23 & B & 5.0 & 5 & 5.0  & 0.98  & 1.020  & 0.94  & 18.3 & 77.9 &   11.07 & \\\\ HD\\,101636 &  & & 0.3454 & 5 & 0.29 & 0.86 & 1.026 & 0.09 & 30.8 & 71.0 &   9.41 & \\\\ \\hline \\end{tabular} \\begin{list} {}{} \\item[${\\boldmath^{\\mathrm{1}}}$] {\\scriptsize the grade of each observing block (OB) is set by the on-site VLT observer based on the user-requested ambient conditions (seeing, sky transparency, airmass) vs. the actual observing conditions. A grade ``A\" means that the conditions were fully within specifications, ``B\" that they were mostly within specifications,  while ``C\" means they were out of specifications and the OB should be repeated if possible. } \\item[${\\boldmath^{\\mathrm{2}}}$] {\\scriptsize on source = DIT $\\times$ NDIT $\\times$ number of OBJECT frames} \\item[${\\boldmath^{\\mathrm{3}}}$] {\\scriptsize average astronomical site monitor seeing at start in the visible (0.5 $\\mu$m) and at the zenith (airmass=1). This value can differ significantly from the NACO image quality, which is usually better. ``NA\" means that no seeing data was available for this night. } \\item[${\\boldmath^{\\mathrm{4}}}$] {\\scriptsize average coherence time of the atmosphere at 0.5 $\\mu$m calculated by the real time computer (RTC) on real data. Short values of  $\\tau_{0}$~($<$~6 ms) represent a seeing not stable in time} \\item[${\\boldmath^{\\mathrm{5}}}$] {\\scriptsize computed with \\emph{eclipse} in the final processed image}% \\item[${\\boldmath^{\\mathrm{6}}}$] {\\scriptsize of the star in the processed image adopting a Moffat fitting to the radial profile. The {\\bf full width at half-maximum (FWHM)} of the diffraction-limited PSF for the VLT at ~2.12~$\\mu$m is 65~mas} \\item[${\\boldmath^{\\mathrm{7}}}$] {\\scriptsize optical magnitude retrieved from SIMBAD: 1) DoAr\\,21 - \\cite{Cieza2007}, 2) HD\\,135344B - \\cite{Hog2000}, 3) HR\\,4796A -  \\cite{Barrado2006}, 4) T\\,Cha - \\cite{Tanner2007}, 5) TW\\,Hya - \\cite{Torres2006}} \\item[${\\boldmath^{\\mathrm{8}}}$] {\\scriptsize DoAr\\,21: PSF saturated in the central pixels (4 to 7 pixel range); T\\,Cha: elliptical along x-axis; PSF calibrators HD\\,144156 and HD\\,94889 have close companions at  d=1\\farcs33, PA=20$^\\circ$ and d=2\\farcs2, PA=260$^\\circ$,~respectively.} \\end{list} \\end{table*}  The transition from the gas/dust-rich actively accreting optically thick disks surrounding T-Tauri and Herbig Ae/Be stars to the gas/dust-poor optically thin debris disks is currently not well understood. Of the several processes proposed for disk evolution and dissipation, the formation of planets remains the most exci\\-ting \\citep[e.g.,][and references therein]{Meyer2007a}. Early stu\\-dies of young stars dating back to \\emph{IRAS} \\citep[e.g., ][]{Strom1989,Skrutskie1990} identified several sources with spectral ener\\-gy distributions (SEDs) that showed a deficit of mid-infrared flux and a rise into the far-IR. These were interpreted as a sign of dust clearing, which is expected from disk evolution as the dust grows and settles or is removed with the gas by viscous accretion or other dispersal mechanisms such as photoevapo\\-ration and tidal truncation because of close companions, both stellar and substellar. More recent surveys of young stellar clusters with {\\it Spitzer} revealed a small population of  these sources in several nearby star-forming regions (d $\\le$ 400 pc) based on photometry \\citep{Sicilia-Aguilar2006,Muzerolle2010}, spectroscopy \\citep{Brown2007,Furlan2009, Oliveira2010} or both \\citep{Sicilia-Aguilar2008,Merin2010}.  They are  currently  referred to as `transitional' or `cold' disks after \\cite{Brown2007},  because of their lack of emission from warm dust. The reported fraction of transitional disks can vary from a few to 50\\% depending on definition and the method of classification. A consensus in nomenclature is still lacking and, as demons\\-trated in \\cite{Merin2010}, a significant fraction of transitional disks classified with photometry turns out to be not transitional when using spectroscopy. The general definition of a transitional disk is a young disk  ($<$~10~Myr) with an optically thin inner region (the ``opacity hole/gap\") surrounded by an optically thick outer disk (beyond the hole/gap outer radius, R$_{\\mathrm {hole/gap,out}}$). The mid-infrared emission originates in the inner edge or  ``wall\" of the truncated disk that is directly illuminated by the star \\citep{Calvet2005}. Gapped disks that still retain an inner component of warm optically thick dust, typi\\-cally of less than a few AU in width, are also referred to as `pre-transitional' after \\cite{Espaillat2007}.  Most transitional disks are still accreting, so that gas is being transported through the inner cleared hole/gap onto the star. In some sources a small amount of micron  or sub-micron dust coexists with the gas in the hole/gap region and causes an excess over photospheric fluxes in the near-IR.  The detection of tidal gaps in some disks has confirmed the inference of the inner holes from unresolved data. These gaps and holes have already been resolved with high-resolution contrast imaging and interferometry in the optical, infrared and at (sub)millimeter wavelengths in a handful of objects:  LkCa\\,15 \\citep{Thalmann2010,Pietu2006}, TW\\,Hya \\citep{Hughes2007}, LkH$\\alpha$\\,330 \\citep{Brown2008,Brown2009}, GM\\,Aur \\citep{Dutrey2008,Hughes2009}, HR\\,4796A \\citep{Schneider1999,Schneider2009}, SR\\,21N, WSB\\,60, DoAr\\,44,  HD\\,135344B \\citep{Andrews2009,Brown2009}, and T\\,Cha \\citep{Olofsson2011}. The direct images and SEDs modeling of transitional disks suggests inner cavities in the dust distribution of  3 to 100 AU in extension \\cite[e.g.,][]{Dalessio2005b,Najita2007,Brown2007,Cieza2010,Olofsson2011}. Although se\\-ve\\-ral mechanisms have been proposed to explain the origin of inner disk holes or cleared gaps,  the exis\\-tence of previous\\-ly unresolved close binaries has been confirmed in se\\-ve\\-ral transitional disks  \\citep{Ireland2008,Pott2010} and is the first scenario to check before making any other assumptions.   This paper presents near-infrared high-spatial resolution narrow-band imaging of five well known transitional disks  --  DoAr\\,21, HD\\,135344B (SAO\\,206462), HR\\,4796A, T\\,Cha, and TW\\,Hya --  and is aimed at detecting close companions within their inner hole/gap region. The VLT/NACO observations and data reduction are presented in \\S~\\ref{observations}. The results, including the PSF-subtracted images and the sensitivity curves are shown in \\S~\\ref{results}, where we also refer to the mass limits and orbit cons\\-traints of possible companions. In \\S~\\ref{discussion} we tentatively classify the five disks on basis of the mechanisms that potentially create their inner opacity hole/gap.  We also ana\\-ly\\-ze the results for each target independently and discuss how the constraints derived from the NACO images complement or support previous studies of companions that made use of other techniques. The conclusions of this study are given in \\S~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  This paper addresses one possible origin for inner holes or  cleared gaps in protoplanetary disks. These holes/gaps  are seen as potential signposts of planet formation but other explanations are possible. Here we investigate the mechanism of tidal truncation by close binaries as the possible or alternate explanation for the origin of the inner holes/gaps observed in five transitional disks. The conclusions are that:  \\begin{enumerate} \\item With the proviso of possible unfavorable projections, the VLT/NACO images rule out the presence of binary compa\\-nions from 0\\farcs1 to $\\sim$~7\\arcsec~in all five transitional disks and down to a 5$\\sigma$ detection limit of  2.2--4.9 mag in contrast for the inner region (at 0\\farcs1) and 7.7--9.3~mag (4.7 mag in the case of T\\,Cha) for the outer region (beyond 1\\farcs0).  \\item  For the transitional disks DoAr\\,21, HD\\,135344B, HR\\,4796A, and T\\,Cha, the VLT/NACO observations resolve part of the inner hole/gap region of the disk and can potentially discard binaries down to a certain mass limit as the main explanation for the inner opacity cavities. For these objects, our results reduce the suite of potential mechanisms that create the holes to closer or lower-mass companions, giant planet formation, efficient grain growth or photoevaporation (for DoAr\\,21 and HR\\,4796A).  \\end{enumerate}"
    },
    "1107/1107.4885_arXiv.txt": {
        "abstract": "We investigate how the conditions occurring in a protoplanetary disc may determine the final structure of a planetary system  emerging from such a disc. We concentrate our attention on the dynamical interactions between disc and planets leading to orbital migration, which in turn, in favourable circumstances, can drive planets into a mean-motion commensurability.  We find that for a system containing a gas giant on the external orbit and a Super-Earth on the internal one, both embedded in a gaseous disc, the 2:1 resonance is a very likely configuration, so one can expect it as an outcome of the early phases of the planetary system formation. Our conclusion is based on an extensive computational survey in which we ask what are the disc properties (the surface density and the viscosity) for which the 2:1 commensurability may be attained. To answer this question we employ a full two-dimensional hydrodynamic treatment of the disc-planet interactions.  In general terms, we can claim that the conditions which favour the 2:1 mean-motion resonance exist in the protoplanetary discs with mass accretion rates in the range of 5$\\cdot 10^{-9} - 5\\cdot 10^{-8}$ M$_{\\odot}$/year. For accretion rates higher than those needed for the 2:1 commensurability we observe a variety of behaviours, among them  the passage to the 3:2 resonance, the scattering of the Super-Earth or the divergent migration caused by the outward migration of the gas giant. The results we have obtained from numerical simulations are compared with the predictions coming from the existing analytical expressions of the migration speed and the strength of the mean motion resonances. The conditions that we have found for the attainment of the 2:1 commensurability are discussed in the framework of the properties of protoplanetary discs that are known from the observations. ",
        "introduction": "\\label{introduction} Mean-motion resonances are ubiquitous phenomena in  planetary systems. They carry  important information about the formation processes and the further evolution of those systems. Moreover, their occurrence may increase the  chance of detecting planets which happen to be in a resonant configuration \\citep{agol2005, holman}. They might also be  able to tell us more about the properties of the protoplanetary discs in which  planets are born. The increasing number of low-mass planets being discovered in  multiple planetary systems \\citep{mayor9a, mayor9b, lovis10} is the first strong observational motivation for  the studies of planetary formation and evolution. Another motivation is connected with the growing amount of data coming from  high angular resolution observations of protoplanetary discs \\citep{Andrews09, hughes, isella, olo, huelamo}. The excitement in this area of research is stimulated also by the fast progress in theoretical investigations of the orbital migration due to the action of tidal torques \\citep{linpap93, ward97, tanaka02, paapap08, paapap09}. The convergent migration of planets or protoplanetary cores is  one of the most promising processes to explain the formation of resonant configurations.  Extensive studies of  possible resonant structures have been performed among others by \\cite{nelpap2002} and \\cite{papszusz2005,papszusz2010}. In  \\cite{nelpap2002} the possible commensurabilities to be expected in the system of planets in the Jovian mass range are investigated,  while in \\cite{papszusz2005,papszusz2010}  migrating planets in the terrestrial mass range are considered and the  conditions for the occurrence of first order mean-motion resonances are provided. \\cite{ fonel07b}, \\cite{zhou}, \\cite{raymond06}, \\cite{mandell} have shown that terrestrial planets can grow and be retained in the hot-Jupiter systems and that their resonant captures by a gas giant are quite common. In the present paper we consider a system containing a Jupiter-like planet evolving together with a Super-Earth in a gaseous disc. We perform a survey of the different disc properties in which this system is embedded using full hydrodynamical calculations to simulate the dynamical interactions between discs and planets. Our purpose here is to find out what are the conditions for the occurrence of the 2:1 commensurability in such systems  and at the same time to get an insight into the environment suitable for the formation of planets on the basis of the existence of certain mean-motion commensurabilities in  young  planetary systems. It has been shown in \\cite{paperI, paperII} that in the gas giant -- Super-Earth systems the planets embedded in a gaseous disc are easily locked in the commensurabilities if and only if the Super-Earth is located inside the orbit of the gas giant. For this reason, we consider here exactly that kind of configurations in which the Super-Earth is the inner and the gas giant is the outer planet of the system. In the previous paper \\citep{paperI}, where the same system has been studied, the 2:1 commensurability was not included because of computational limitations.  This paper is organized as follows: In Section \\ref{survey} we describe how our computational survey has been performed. Next, we illustrate the relevant features of the evolution of Jupiter-like planets in discs with different viscosities. In Section \\ref{results} we give the results of our survey presenting different configurations achieved during the tidal evolution of the planets in the disc.  These results are compared in Section \\ref{theoretical} with the theoretical expectation based on the analytical and semi-analytical expressions for the migration speed of the planets and the strength of the mean motion resonances. In Section \\ref{observational}  we point those observed protoplanetary discs in which the physical conditions that favour the occurrence of the 2:1 resonance are present. Finally, our findings are summarized and discussed in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} In this paper we have constrained the properties of protoplanetary discs which favour the locking of a Super-Earth and a Jupiter-like planet orbiting around a Sun-like star in the 2:1 commensurability. The Super-Earth and the Jupiter are  located on the internal and external orbits respectively. The results are illustrated in Figs. \\ref{fig1} and \\ref{fig10}. In both figures the medium grey region is where we expect the 2:1 resonance to occur. In the low surface density regime (less or equal to $4\\cdot 10^{-4}$) the final outcome is relatively easy to determine. The migration of both planets is rather slow in this region. We expect the 2:1 resonance for viscosities not lower than $4\\cdot 10^{-6}$. Below this threshold the Jupiter migrates slower than the Super-Earth and no commensurability can be achieved. On the other side, in order to get the 2:1 resonance, the viscosity cannot be too high, because then the Jupiter will migrate outwards. This is the case in which the viscosity combined with the surface density through Eq. (\\ref{signu}) gives accretion rates higher than $10^{-7}$ M$_{\\odot}$/yr. In the high surface density regime (higher than $6\\cdot 10^{-4}$) the situation is more complicated  and the 2:1 resonance is allowed only in a narrow strip where the viscosity is varying in the range from $4 \\cdot 10^{-6}$ till $10^{-5}$. In this regime  the equilibrium gap profile is not empty. This is particularly true for relatively high viscosities. For that reason, putting the gas giant in the disc results in  its very rapid migration. To avoid the artefact caused by the process of gap opening, we kept the Jupiter-like planet fixed, not allowing for its migration till the gap has reached the quasi-stationary state and stops changing significantly. In this way we have obtained the 2:1 resonance for viscosities ranging from $4\\cdot 10^{-6}$ till $10^{-5}$. The results coming from numerical simulations have been compared with the analytical predictions summarized in Fig. \\ref{fig11} and with the current observations, see Fig. \\ref{fig10}.   The orbital migration of planets embedded in  protoplanetary discs is a natural mechanism in which  resonant configurations might occur. However, it is not the only way to form  orbital commensurabilities. It has been suggested that the planet-planet scattering might be an efficient mechanism leading to the resonant capture of planets \\citep{raymond}. Our computational set up allows us to investigate also this scenario in the situation when the gaseous disc is still present in the system. We have employed a simple procedure to locate both planets in a disc with high surface density (higher than $6\\cdot 10^{-4}$) and let them migrate without waiting for the Jupiter to open a gap. Here we discuss  in details one example of resonance obtained in this way. Our probe, which consists of a Jupiter-like planet and a Super-Earth, is located in a disc with  surface density equal to $9\\cdot 10^{-4}$ and  viscosity value taken to be 8$\\cdot 10^{-6}$. As expected, the Jupiter migrates extremely fast  (with type III migration, \\cite{massetpap}), traveling one third of the distance from its initial location to the star in just 100 orbits. Its orbit crosses the orbit of the slowly migrating Super-Earth and we witness a resonant capture into 1:1 mean motion resonance. This is shown in Fig. \\ref{fig9} (two upper panels). After the occurrence of the locking the Super-Earth stays very close to the Jupiter on the horseshoe orbit with an eccentricity of about 0.03 (Fig. \\ref{fig9}, left lower panel). The resonant angle is shown in Fig. \\ref{fig9} (right lower panel). The 1:1 commensurability created in the way described above lasted for roughly 700 orbits and after that the Super-Earth has been ejected from the system.  Concluding, we have shown how  the mean-motion resonances observed in the already formed planetary systems can help in getting additional insight into the conditions occurring in protoplanetary discs in the early phases of the planetary system evolution. In other words, we  have made an attempt to  constrain the resonant properties of the  planetary systems which may occur in observed protoplanetary discs. Assuming that the resonances are due to early stages of the evolution, we can answer the question on what should be the disc properties in order to attain  the 2:1 resonance."
    },
    "1107/1107.4627_arXiv.txt": {
        "abstract": "In this work, we obtain isotropic extensions of the usual spherically symmetric vacuum geometries in general relativity. Exact and perturbative solutions are derived. The classes of geometries obtained include black holes in compact and noncompact universes, wormholes in the interior region of cosmological horizons, and anti-de Sitter geometries with excess/deficit solid angle. The tools developed here are applicable in more general contexts. ",
        "introduction": "Spacetimes described by spherically symmetric solutions of Einstein's equations are of paramount importance both in astrophysical applications and theoretical considerations. And among those, black holes are highlighted. They are relevant sources of gravitational radiation, offering possible observational signatures of general relativity extensions. The current and upcoming gravitational wave experiments and the possibility of detecting black holes in accelerators are strong motivations for the investigation of such models.  In a vacuum, Birkhoff's theorem and its generalizations to nonasymptotically flat cases uniquely fix the metric as the Schwarzschild, Schwarzschild--de Sitter or Schwarzschild--anti-de Sitter geometries, the vacuum solutions of the usual general relativity with zero, positive or negative values for the cosmological constant, respectively. Nevertheless, our universe is not in a vacuum state, even if its dynamics could principally be driven by an (at least) effective cosmological constant \\cite{Riess,Perlmutter}. Therefore, it is interesting to consider how the compact solutions in general relativity are modified in a cosmological scenario. In a different direction, anti-de Sitter geometries gained interest due to the proposed anti-de Sitter (AdS)--conformal field theory (CFT) correspondence \\cite{Maldacena,Witten}. This conjecture proposes a duality between gravity in AdS spaces and CFTs. A better understanding of AdS backgrounds plays an important role in this program.  In this work we are mainly interested in black holes in a cosmological environment. Of the two main assumptions of the cosmological principle, homogeneity is lost when compact objects are considered. Nevertheless isotropy is still possible, and we enforce this condition. Within this context, we investigate spatially isotropic solutions close -- continuously deformable -- to the usual vacuum solutions.  From a more mathematical point of view, isotropy is a condition that can be implemented in a purely coordinate invariant way, and as a simple linear constraint in the usual coordinate system $(t,r,\\theta,\\phi)$. This geometrical condition is somewhat akin to the imposition of a constant Ricci scalar, natural in a Randall--Sundrum brane world context \\cite{Randall-Sundrum,Shiromizu}. Solutions of this constraint can also be expressed as a linear deformation of the general relativity vacuum solutions, whose extensions were treated in \\cite{Casadio,Bronnikov2,Bronnikov,Molina1}.  The present work deals with the problem of extending the usual Einstein's equation solutions in a similar manner. While one practical approach to general relativity is to specify the matter content and from that the spacetime structure, this is not the only possible treatment. For instance, one can specify the spacetime metric based on physical and geometrical considerations, and later follow its implications for the energy-momentum tensor \\cite{Casadio,Bronnikov2,Bronnikov,Molina1,Bronnikov:2008,GonzalezDiaz:2011,Morris:1988cz,Morris:1988PRL,Jacobson:2007}. In fact, currently the cosmological scientific community is attempting to adjust the energy density and pressure needed to produce an accelerated FLRW cosmology. We favor here this latter approach.  The structure of this paper is presented in the following. In Sec. \\ref{secII} we develop the basic formalism to be used. In Sec. \\ref{secIII} we apply this formalism deriving linear solutions which exactly satisfy the isotropy constraint. A large class of structures appear in the process. In Sec. \\ref{secIV}, we go beyond the linear cases, obtaining more general backgrounds which are approximately isotropic. Some final comments are made in Sec. \\ref{secV}. ",
        "conclusions": "\\label{secV}  In this work we have shown that the vacuum spherically symmetric geometries, Minkowski, Schwarzschild, de Sitter, Schwarzschild--de Sitter, anti-de Sitter and Schwarzschild--anti-de Sitter, can be isotropically deformed to take into account the existence of some material content. Even considering linear deformations, in the sense that the physical quantity $\\left[T_{\\nu}^{\\mu}\\right]$ is strictly linear in $\\delta$, we have obtained a zoo of geometries containing usual or exotic astronomical objects with different asymptotic behaviors.  In particular, when considering linear deformations of the Minkowski solution ($\\Lambda=0$ and $M=0$), we have obtained spatially homogeneous and isotropic universes which are spatially closed (Einstein universe) and open, for negative and positive values of the deformation parameter ($\\delta$, which in Sec. \\ref{secIII} is included in $D$, with ${\\rm sign}\\left(\\delta\\right)={\\rm sign}\\left(D\\right)$ and, in Sec. \\ref{A1}, ${\\rm sign}\\left(\\delta\\right)\\neq{\\rm sign}\\left(K\\right)$), respectively. It is well known that, usually in order to consider those models, a cosmological constant is needed. Nevertheless, in this case the cosmological constant is not appearing through the Einstein equations, but it is part of the fluid related to the deformation.  The deformation of a Schwarzschild geometry could lead to a spacetime where a black hole is in a noncompact universe or two black holes, the original one and the reflected one, in a compact universe, depending on the sign of $C$ (which is the same that the sign of $\\delta$). About this second solution, it must be pointed out that we have been able to obtain a closed structure with two black holes, because the deformation implies that we are no longer considering a vacuum background where Birkhoff's theorem holds (as studied in Ref.~\\cite{Uzan:2010nw}).  We have also shown that both de Sitter and Schwarzschild--de Sitter spacetimes can be smoothly deformed into a geometry which can be interpreted to describe a wormhole-like structure in an asymptotically de Sitter universe if $C<0$. Such a wormhole would consistently be supported by a material content, originated by the deformation, which violates the null energy condition on and around its throat.  The deformation of the anti-de Sitter geometry leads to a spacetime which asymptotically behaves as an anti-de Sitter space with a deficit or excess of solid angle, or a compact static universe, depending on the value of $C$. A black hole should be considered in those geometries when one is deforming a Schwarzschild--anti-de Sitter space.  The above mentioned spacetimes which show the presence of black holes are examples of structures in nonvacuum spacetimes which are in equilibrium with their environment. We have noted that in all the studied cases the fluid originated by the deformation is such that its equation-of-state parameter is equal to minus one in the black hole Killing horizon. Therefore, we have hypothesized a nonaccretion condition, in some sense inspired in Refs.~\\cite{Babichev:2004yx,MartinMoruno:2008vy}, which state that a fluid which behaves as a cosmological constant on the black hole horizon would not be accreted by it. It must be emphasized that this seems to be a sufficient but not necessary condition in order to have no accretion, at least in principle.  All the geometries obtained consistently reduce to the vacuum cases when $C\\rightarrow 0$. The material content originated by the deformation seems to have complicated functional forms in some of the considered backgrounds. Nevertheless, as we have seen in some cases, the energy densities and pressures can be interpreted as resulting from the sum of two or more fluids with simpler forms.  On the other hand, we have also briefly considered the application of the deformation formalism relaxing the strict isotropy constraint. As we have shown, in this case additional specifications are necessary to obtain unique solutions. We have considered possible choices of the perturbation based on physical grounds. Nevertheless, it must be pointed out that this is a powerful formalism which could provide us with new interesting solutions.  Moreover, although isotropy is the key point in the present work, extensions of usual solutions in general relativity subjected to other constraints can be explored within the same approach. For example, a constraint equation in the form $p_{r}=w_{r}\\,\\rho$ with a constant value of $w_{r}$ admits as an exact linear solution (with $\\Lambda = 0$) \\begin{equation} A(r) = 1 - \\frac{2M}{r}  \\,\\, , \\end{equation} \\begin{equation} B(r) = 1 - \\frac{2M}{r} + \\frac{\\delta}{r} \\left( r - 2M \\right)^{-1/w_{r}} \\,\\, . \\end{equation} A rich set of structures can be obtained for the different values of $w_{r}$ taken. This constraint could be relevant in physical scenarios."
    },
    "1107/1107.0909_arXiv.txt": {
        "abstract": "{} {We present a unique multi-epoch infrared interferometric study of the oxygen-rich Mira variable RR~Aql in comparison to radiative transfer models of the dust shell. We investigate flux and visibility spectra at 8 -- 13\\,$\\mu$m with the aim of better understanding the pulsation mechanism and its connection to the dust condensation sequence and mass-loss process.} {We obtained 13 epochs of mid-infrared interferometry with the MIDI instrument at the VLTI between April 2004 and July 2007, covering minimum to pre-maximum pulsation phases (0.45--0.85) within 4 cycles. The data are modeled with a radiative transfer model of the dust shell where the central stellar intensity profile is described by a series of dust-free dynamic model atmospheres based on self-excited pulsation models. We examined two dust species, silicate and Al$_2$O$_3$ grains. We performed model simulations using variations in model phase and dust shell parameters to investigate the expected variability of our mid-infrared photometric and interferometric data. } {The observed visibility spectra do not show any indication of variations as a function of pulsation phase and cycle. The observed photometry spectra may indicate intracycle and cycle-to-cycle variations at the level of 1--2 standard deviations. The photometric and visibility spectra of RR~Aql can be described well by the radiative transfer model of the dust shell that% uses a dynamic model atmosphere describing the central source. The best-fitting model for our average pulsation phase of $\\overline{\\Phi_V}=0.64\\pm0.15$ includes the dynamic model atmosphere M21n ($T_\\mathrm{model}=2550$\\,K) with a photospheric angular diameter of $\\theta_\\mathrm{Phot}=7.6\\pm0.6$\\,mas, and a silicate dust shell with an optical depth of $\\tau_V=2.8\\pm0.8$, an inner radius of $R_\\mathrm{in}=4.1\\pm0.7\\,R_\\mathrm{Phot}$, and a power-law index of the density distribution of $p=2.6\\pm0.3$. The addition of an Al$_2$O$_3$ dust shell did not improve the model fit. However, our model simulations indicate that the presence of an inner Al$_2$O$_3$ dust shell with lower optical depth than for the silicate dust shell can not be excluded. The photospheric angular diameter corresponds to a radius of $R_\\mathrm{phot}=520^{+230}_{-140}R_\\odot$ and an effective temperature of $T_\\mathrm{eff}\\sim2420\\pm200$\\,K. Our modeling simulations confirm that significant intracycle and cycle-to-cycle visibility variations are not expected for RR~Aql at mid-infrared wavelengths within our uncertainties.} {We conclude that our RR~Aql data can be described by a pulsating atmosphere surrounded by a silicate dust shell. The effects of the pulsation on the mid-infrared flux and visibility values are expected to be less than about 25\\% and 20\\%, respectively, and are too low to be detected within our measurement uncertainties. } ",
        "introduction": "Asymptotic giant branch (AGB) stars are low-to-intermediate mass stars at the end of their stellar evolution. These stars, including Mira type stars, exhibit many complex processes such as shock fronts propagating through the stellar atmosphere, large-amplitude pulsation, and molecule and dust formation leading to strong mass loss via a dense and dusty outflow from an extended stellar atmosphere \\citep{Andersen2003}. The mass loss rates can reach up to $10^{-4}M_\\odot$yr$^{-1}$ \\citep{Matsuura2009} with expansion velocities of 5-30 km s$^{-1}$ \\citep{Hofner2007}. Dust grains are created by condensation from the gas phase in the very cool and dense outer part of the pulsating atmosphere. Dust created by carbon-rich stars owing to its high opacity absorbs radiation from the star, and the wind is believed to be driven by radiation pressure on the dust particles that drag gas along via collisions (H{\\\"o}fner 2008).  Despite remarkable progress in theoretical and observational studies, it still remains unclear whether the dust in oxygen-rich envelopes is opaque enough to drive the observed massive outflows \\citep{Woitke2006,Hofner2008}. Because of the wind, all the matter around the core of the star is eventually returned to the interstellar medium, and the star evolves toward the planetary nebula phase, leaving the core as a white dwarf. Generally many aspects of the physics of AGB stars remain poorly understood, such as the detailed atmospheric stratification and composition of the stars, or the role of stellar pulsation and its connection to the dust formation and massive outflow. Improving our understanding of the physical processes leading to the substantial mass loss is important, as AGB stars play a crucial role in the chemical enrichment and evolution of galaxies by returning gas and dust to the interstellar medium (ISM).  Thanks to their large diameters and high luminosities, Mira variables are ideal targets for high angular resolution observations. Near-infrared interferometric (NIR) techniques provide detailed information regarding the conditions near the continuum-forming photosphere such as effective temperature, center-to-limb intensity variations, the stellar photospheric diameter, and its dependence on wavelength and pulsation phase \\citep[e.g.][]{Haniff1995,Perrin1999,vanBelle1996,Young2000a,Thompson2002a,Thompson2002b,Quirrenbach1992,Fedele2005,Ohnaka2004,Millan-Gabet2005,Woodruff2008}. New theoretical self-excited dynamic model atmospheres of oxygen-rich stars have been created and successfully applied \\citep{Tej2003,Hofmann1998,Ireland2004a,Ireland2004b,Ireland2008,Wittkowski2008}.  Observations at mid-infrared wavelengths are well-suited to studying the molecular shells and the dust formation zone of evolved stars. Mid-infrared interferometry is sensitive to the chemical composition and geometry of dust shells, their temperature, inner radii, radial distribution, and the mass loss rate \\citep[e.g.][]{Danchi1994,Monnier1997,Monnier2000,Weiner2006}. \\citet{Lopez1997} presented long-term observations at 11\\,$\\mu$m of o Ceti obtained with the Infrared Spatial Interferometer (ISI). The observed visibilities change from one epoch to the next and are not consistent with simple heating or cooling of the dust with change in luminosity as a function of stellar phase, but rather with large temporal variations in the density of the dust shell. The data were compared to axially symmetric radiative transfer models and suggest inhomogeneities or clumps. \\citet{Tevousjan2004} has studied the spatial distribution of dust around four late-type stars with ISI at 11.15\\,$\\mu$m, and find that the visibility curves change with the pulsation phase of the star. The dust grains were modeled as a mixture of silicates and graphite. The results suggest that the dust shells appear to be closer to the star at minimum pulsation phase, and farther away at maximum phase, which is also demonstrated by \\citet{Wittkowski2007}. \\citet{Ohnaka2007} reported on temporal visibility variations of the carbon-rich Mira variable V Oph with the instruments VINCI and MIDI at the VLTI. This temporal variation of the N-band angular sizes is largely governed by the variations in the opacity and the geometrical extension of the molecular layers (C$_2$H$_2$ and HCN) and the dust shell (amC + SiC).  Additional information to the near- and mid-infrared observations can be obtained with Very Long Baseline Array (VLBA) observations, which allow one to study the properties of the circumstellar environment of evolved stars using the maser radiation emitted by some molecules, most commonly SiO, H$_2$O, and OH. By observing the maser radiation at tens to several hundred AU from the star, it is possible to determine accurate proper motions and parallaxes \\citep{Vlemmings2003}, total intensity and linear polarizations \\citep{Cotton2008}, and the structure of the environment of the stars \\citep{Diamond1994}.  This work is part of the ongoing project of concurrent multiwavelength observations using optical long baseline interferometry (AMBER and MIDI at the VLTI) combined with radio interferometry (VLBA). We aim to investigate layers at different depths of the atmosphere and circumstellar envelope (CSE) of the stars. Previous results from this project of joint VLTI/VLBA observations for the Mira star S Ori can be found in \\citet{Boboltz2005} and \\citet{Wittkowski2007}. Here we present results of a long-term VLTI/MIDI monitoring of the oxygen-rich Mira variable RR~Aql. In addition to the VLTI/MIDI observation presented in this paper, we obtained coordinated multi-epoch VLBA observations of SiO and H$_2$O masers towards RR~Aql, which will be presented in a subsequent paper.  \\begin{figure} \\includegraphics[height=0.35\\textheight,angle=90]{light_curve_RRAql_AAVSO_AFOEV_proceedings2.ps} \\caption{Visual light curve of RR~Aql based on data from the AAVSO and AFOEV databases as a function of Julian Date and stellar cycle/phase. The green arrows indicate the dates of our VLTI/MIDI observations. The blue arrows indicate the dates of our VLBA observations. Here, the full arrows denote observations of SiO maser emission and simple arrows observations of H$_2$O maser emission. } \\label{LightCurve} \\end{figure} \\begin{table*}[tbp] \\caption[]{VLTI/MIDI observation of RR~Aql.} \\label{observations} \\centering \\begin{tabular}{lrrrrlrllrrrr} \\hline \\hline Ep.&date&Time&JD&$\\Phi$$_{vis}$&Conf.&$B$& Disp.&BC&$B_p$ &P.A.&Seeing&$\\tau$$_0$\\\\ && [UTC]&&&&&Elem.&&[m]&[deg]&[$''$]&[msec]\\\\ &DD/MM/YYYY&&&&&&&&&&&\\\\\\hline A & 09/04/2004 & 07:26 & 2453105 & 0.58 & U2-U3 & 47m & Prism & HS & 33.57 & 3.21 & 0.45 & 7.4 \\\\ A & 10/04/2004 & 09:05 & 2453106 & 0.58 & U2-U3 & 47m & Prism & HS & 37.17 & 26.09 & 0.76 & 7.8 \\\\ A &\t     & 09:52 & 2453106 & 0.58 & U2-U3 & 47m & Prism & HS & 39.86 & 33.48 & 0.64 & 8.2 \\\\ B & 09/07/2004 & 08:25 & 2453196 & 0.81 & U2-U3 & 47m & Prism & HS & 45.23 & 44.83 & 0.47 & 4.8 \\\\ C & 28/07/2004 & 08:09 & 2453215 & 0.86 & U2-U3 & 47m & Prism & HS & 42.54 & 41.57 & 1.36 & 1.0 \\\\ C & 29/07/2004 & 02:05 & 2453216 & 0.86 & U2-U3 & 47m & Prism & HS & 37.83 & 28.16 & 1.66 & 1.0 \\\\ C &\t     & 03:33 & 2453216 & 0.86 & U2-U3 & 47m & Prism & HS & 42.84 & 39.49 & 0.75 & 2.3 \\\\ C &          & 04:55 & 2453216 & 0.86 & U2-U3 & 47m & Prism & HS & 45.99 & 44.67 & 1.04 & 1.6 \\\\ C &          & 06:33 & 2453216 & 0.86 & U2-U3 & 47m & Prism & HS & 46.16 & 45.71 & 3.38 & 0.5 \\\\ C &          & 07:04 & 2453216 & 0.86 & U2-U3 & 47m & Prism & HS & 45.32 & 44.92 & 2.68 & 0.7 \\\\ C & 01/08/2004 & 00:12 & 2453219 & 0.87 & U2-U3 & 47m & Prism & HS & 33.76 & 6.86 & 0.73 & 3.1 \\\\ D & 18/04/2005 & 10:30 & 2453479 & 1.53 & U2-U4 & 89m & Prism & HS & 88.37 & 81.72 & 0.76 & 3.5 \\\\ D & 19/04/2005 & 07:53 & 2453480 & 1.53 & U2-U4 & 89m & Prism & HS & 61.36 & 76.41 & 0.83 & 2.3 \\\\ E & 20/07/2005 & 03:51 & 2453572 & 1.76 & U1-U4 & 130m & Prism & HS & 119.63 & 59.78 & 0.71 & 1.9 \\\\ E &          & 06:16 & 2453572 & 1.76 & U1-U4 & 130m & Prism & HS & 128.64 & 63.28 & 0.81 & 1.6 \\\\ E & 22/07/2005 & 05:10 & 2453574 & 1.77 & U2-U3 & 47m & Prism & HS & 45.65 & 44.12 & 1.00 & 1.4 \\\\ E &          & 07:56 & 2453574 & 1.77 & U2-U3 & 47m & Prism & HS & 44.44 & 43.96 & 1.11 & 1.2 \\\\ F & 18/04/2006 & 09:38 & 2453844 & 2.45 & D0-G0 & 32m & Prism & HS & 28.99 & 70.01 & 1.08 & 2.1 \\\\ F & 19/04/2006 & 08:32 & 2453845 & 2.46 & D0-G0 & 32m & Prism & HS & 24.74 & 65.81 & 0.80 & 6.0 \\\\ F &          & 09:21 & 2453845 & 2.46 & D0-G0 & 32m & Prism & HS & 28.21 & 69.32 & 0.78 & 6.3 \\\\ G & 21/05/2006 & 05:53 & 2453877 & 2.54 & E0-G0 & 16m & Prism & HS & 10.94 & 62.10 & 0.56 & 4.0 \\\\ G & 25/05/2006 & 06:35 & 2453881 & 2.55 & A0-G0 & 64m & Prism & HS & 53.17 & 67.77 & 0.55 & 5.1 \\\\ G &          & 08:46 & 2453881 & 2.55 & A0-G0 & 64m & Prism & HS & 63.93 & 72.67 & 0.44 & 6.7 \\\\ H & 18/06/2006 & 03:58 & 2453905 & 2.61 & D0-G0 & 32m & Prism & HS & 21.37 & 61.34 & 1.29 & 1.5 \\\\ H &          & 04:41 & 2453905 & 2.61 & D0-G0 & 32m & Prism & HS & 25.09 & 66.21 & 1.27 & 1.6 \\\\ H & 18/06/2006 & 05:39 & 2453905 & 2.61 & D0-G0 & 32m & Prism & HS & 29.05 & 70.06 & 1.00 & 2.1 \\\\ H & 20/06/2006 & 05:49 & 2453907 & 2.61 & A0-G0 & 64m & Prism & HS & 59.89 & 70.81 & 1.05 & 3.9 \\\\ H &          & 06:33 & 2453907 & 2.61 & A0-G0 & 64m & Prism & HS & 63.00 & 72.15 & 0.88 & 4.5 \\\\ H & 21/06/2006 & 07:08 & 2453908 & 2.62 & E0-G0 & 16m & Prism & HS & 16.00 & 72.76 & 1.01 & 2.2 \\\\ H &          & 09:36 & 2453908 & 2.62 & E0-G0 & 16m & Prism & HS & 13.21 & 70.37 & 0.77 & 2.9 \\\\ H & 23/06/2006 & 03:20 & 2453910 & 2.62 & E0-G0 & 16m & Prism & HS & 9.84 & 58.42 & 0.76 & 3.9 \\\\ H &          & 04:07 & 2453910 & 2.62 & E0-G0 & 16m & Prism & HS & 11.97 & 64.86 & 0.67 & 4.4 \\\\ I & 08/08/2006 & 05:12 & 2453956 & 2.74 & A0-G0 & 64m & Prism & HS & 61.41 & 72.62 & 1.02 & 2.3 \\\\ I &          & 05:57 & 2453956 & 2.74 & A0-G0 & 64m & Prism & HS & 56.88 & 71.58 & 1.10 & 2.2 \\\\ I & 09/08/2006 & 06:39 & 2453957 & 2.74 & A0-G0 & 64m & Prism & HS & 50.33 & 69.49 & 1.47 & 1.5 \\\\ I & 10/08/2006 & 03:07 & 2453958 & 2.74 & D0-G0 & 32m & Prism & HS & 31.37 & 72.03 & 1.43 & 2.6 \\\\ I &          & 04:31 & 2453958 & 2.74 & D0-G0 & 32m & Prism & HS & 31.65 & 72.90 & 1.67 & 2.5 \\\\ I &          & 05:45 & 2453958 & 2.74 & D0-G0 & 32m & Prism & HS & 28.72 & 71.74 & 1.43 & 3.2 \\\\ I & 11/08/2006 & 02:44 & 2453959 & 2.74 & D0-G0 & 32m & Prism & HS & 30.76 & 71.51 & 1.43 & 2.3 \\\\ I & 13/08/2006 & 04:05 & 2453961 & 2.75 & E0-G0 & 16m & Prism & HS & 15.94 & 72.90 & 0.91 & 1.9 \\\\ I &          & 04:56 & 2453961 & 2.75 & E0-G0 & 16m & Prism & HS & 15.26 & 72.56 & 0.63 & 2.7 \\\\ I &          & 05:52 & 2453961 & 2.75 & E0-G0 & 16m & Prism & HS & 13.76 & 71.07 & 0.60 & 2.8 \\\\ J & 14/09/2006 & 01:01 & 2453993 & 2.83 & E0-G0 & 16m & Prism & HS & 15.82 & 72.29 & 0.85 & 2.3 \\\\ J & 16/09/2006 & 02:15 & 2453995 & 2.84 & A0-G0 & 64m & Prism & HS & 62.86 & 72.85 & 0.95 & 1.6 \\\\ J &          & 03:13 & 2453995 & 2.84 & A0-G0 & 64m & Prism & HS & 58.12 & 71.90 & 1.68 & 0.9 \\\\ K & 09/05/2007 & 08:53 & 2454230 & 3.43 & D0-H0 & 64m & Prism & HS & 61.43 & 71.46 & 1.12 & 1.3 \\\\ L & 22/06/2007 & 03:50 & 2454274 & 3.54 & G0-H0 & 32m & Prism & HS & 21.96 & 62.20 & 1.17 & 1.0 \\\\ L &          & 05:29 & 2454274 & 3.54 & G0-H0 & 32m & Prism & HS & 29.30 & 70.26 & 1.21 & 0.9 \\\\ M & 03/07/2007 & 07:25 & 2454285 & 3.57 & E0-G0 & 16m & Prism & HS & 15.51 & 72.73 & 1.07 & 2.8 \\\\ M &          & 07:37 & 2454285 & 3.57 & E0-G0 & 16m & Prism & HS & 15.31 & 72.59 & 0.75 & 3.9 \\\\ M & 04/07/2007 & 04:42 & 2454286 & 3.57 & G0-H0 & 32m & Prism & HS & 29.28 & 70.24 & 1.59 & 1.0 \\\\ M &          & 06:30 & 2454286 & 3.57 & G0-H0 & 32m & Prism & HS & 31.98 & 72.83 & 1.51 & 1.1 \\\\ \\hline \\end{tabular} \\tablefoot{The table lists the epoch, the date, the time, the Julian Date (JD), the visual pulsation phase $\\Phi$$_{vis}$, the baseline configuration, the ground length of the configuration, the dispersive element, the beam combiner BC, the projected baseline length $B_p$, the position angle on the sky P.A. (deg. east of north), the DIMM seeing (at 500 nm), and the coherence time $\\tau$$_0$ (at 500 nm).}  \\end{table*} ",
        "conclusions": "We have investigated the circumstellar dust shell and characteristics of the atmosphere of the oxygen-rich Mira variable RR~Aql using mid-infrared interferometric observations. We observed RR~Aql with the VLTI/MIDI instrument at different pulsation phases in order to monitor the photometry and visibility spectra. A total of 57 observations were combined into 13 epochs covering four pulsation cycles between April 2004 and July 2007, and covering pulsation phases between minimum and pre-maximum phases (0.45--0.85).  We modeled the observed data with an ad-hoc radiative transfer model of the dust shell using the radiative transfer code mcsim\\_mpi by \\citet{Ohnaka2006}. In this way, we used a series of dust-free dynamic model atmospheres based on self-excited pulsation models \\citep[M series,][]{Ireland2004a,Ireland2004b} to describe the intensity profile of the central source. This study represents the first comparison between interferometric observations and theoretical models over an extended range of pulsation phases covering several cycles.  Our main observational results are as follows \\begin{itemize} \\item The interferometric data do not show any evidence of intracycle visibility variations. \\end{itemize}  \\begin{itemize} \\item The data do not show any evidence of cycle-to-cycle visibility variations. \\end{itemize}  \\begin{itemize} \\item The 8--13\\,$\\mu$m flux suggests intracycle and cycle-to-cycle photometry variations at a significance level of 1--2$\\sigma$. Follow-up observations with higher accuracy using a dedicated photometric instrument, such as VISIR at the VLT, are needed to confirm this result.\\end{itemize}   These observational results can be explained by dynamic model atmospheres and variations in the dust shell parameters. Simulations using different phases of the dynamic model atmosphere and different sets of dust shell parameters predict visibility variations that are lower than or close to the observed visibility uncertainties (5-20\\,\\%). Model-predicted variations of the photometry spectra are largest around wavelengths of 10\\,$\\mu$m with difference of up to $\\sim$25\\% corresponding to up to 1--2$\\sigma$.  The best-fitting model for our average pulsation phase of $\\overline{\\Phi_V}=0.64\\pm0.15$ includes a silicate dust shell with an optical depth of $\\tau_V (silicate)=2.8\\pm0.8$, an inner radius of $R_\\mathrm{in}=4.1\\pm0.7$\\,$R_\\mathrm{Phot}$, and a power-law index of the density distribution of $p=2.6\\pm0.3$. The corresponding best-fitting atmosphere model of the series used to describe the central intensity profile is M21n ($T_\\mathrm{model}=2550$\\,K, $\\Phi_\\mathrm{model}=0.1$) with a photospheric angular diameter of $\\theta_\\mathrm{Phot}=7.6\\pm0.6$\\,mas. The photospheric angular diameter corresponds to a photospheric radius of $R_\\mathrm{phot}=520^{+230}_{-140}\\,R_\\odot$ and an effective temperature of $T_\\mathrm{eff}\\,\\sim\\,2420\\pm200$\\,K. The latter value is consistent with the effective temperature of the used model M21n. The combined model can reproduce the shape and features of the observed photometry and visibility spectra of RR~Aql well, as well as the SED at 1--40\\,$\\mu$m.      We conclude that our RR~Aql data can be described by a silicate dust shell surrounding a pulsating atmosphere, consistent with observations of Mira variables \\citep{Lorenz-Martins2000}. The effects of the pulsation on the mid-infrared flux and visibility values are expected to be less than about 25\\% and 20\\%, respectively, and are too low to be detected within our measurement uncertainties. Although the addition of an Al$_2$O$_3$ dust shell did not improve the model fit, our simulations also indicate that we cannot exclude the presence of an inner Al$_2$O$_3$ dust shell with relatively low optical depth, which may be an important contributor to the dust condensation sequence.  Our modeling attempt at a radiative transfer model of the dust shell surrounding a dynamic dust-free model atmosphere provides constraints on the geometric extensions of atmospheric molecular and dust shells. It should be noted that it cannot explain the mechanism by which the observed mass loss is produced and the wind is driven. Newer models of M type, i.e. of oxygen-rich Miras, that include atmospheric dust \\citep[e.g. ][Ireland, in preparation]{Ireland2008} typically predict effective acceleration very close to zero in high layers but, so far, no clear outward acceleration, i.e. no wind. Subtle details are being discussed that may overcome this shortcoming of dynamic models of M type Miras \\citep[e.g.][]{Hofner2011}. Future observations aiming at characterizing and constraining other new models of the mass-loss process and the wind driving mechanism at mid-infrared wavelengths would benefit from obtaining more precise photometry values using a dedicated instrument like VISIR at the VLT, the addition of shorter baselines to characterize the extension of the silicate dust shell, and a more complete coverage of the pulsation cycle. The addition of concurrent, spectrally resolved, near-infrared interferometry would be needed to more strongly constrain atmospheric molecular layers located close to the photosphere."
    },
    "1107/1107.5001_arXiv.txt": {
        "abstract": "In this paper we report the identification of two new Galactic O2\\,If*/WN6 stars (WR20aa and WR20c), in the outskirt of the massive young stellar cluster Westerlund 2. The morphological similarity between the near-infrared spectra of the new stars with that of WR20a and WR21a (two of the most massive binaries known to date) is remarkable, indicating that probably they are also very massive stars. New optical spectroscopic observations of WR20aa suggest an intermediate O2 If*/WN6 spectral type. Based on a mosaic made from the 3.6$\\mu$m Spitzer IRAC images of the region including part of the RCW49 complex, we studied the spatial location of the new emission line stars, finding that WR20aa and WR20c are well displaced from the centre of Westerlund 2, being placed at $\\approx 36$ pc (15.7 arcmin) and $\\approx 58$ pc (25.0 arcmin) respectively, for an assumed heliocentric distance of 8 kpc. Also very remarkably, a radius vector connecting both stars would intercept the Westerlund 2 cluster exactly at the place where its stellar density reaches a maximum. We consequently postulate a scenario in which WR20aa and WR20c had a common origin somewhere in the cluster core, being ejected from their birthplace by dynamical interacion with some other very massive objects, perhaps during some earlier stage of the cluster evolution. ",
        "introduction": "\\begin{table*} \\begin{center} \\caption{Coordinates and Optical/NIR photometry of the newly-identified O2\\,If*/WN6 stars in the outskirt of Westerlund 2.\\label{tbl-2}} \\renewcommand{\\thefootnote}{\\thempfootnote} \\begin{tabular}{ccccccccc} \\\\ \\hline\\hline Star & RA     &  Dec & $B$  & $V$ & $J$ & $H$ & $Ks$ & Other designation \\\\ ~ &(J2000) & (J2000) &~    & \\\\ \\hline WR20aa          &10:23:23.49  & -58:00:20.8  &   13.86$\\pm{0.02}$  &   12.69$\\pm{0.02}$  & 9.28$\\pm{0.02}$ & 8.73$\\pm{0.02}$ & 8.40$\\pm{0.02}$ & SS215 \\\\ WR20c           &10:25:02.60  & -57:21:47.3  &   20.10$\\pm{0.05}$  &   17.51$\\pm{0.03}$  & 10.51$\\pm{0.02}$ & 9.57$\\pm{0.02}$ & 9.04$\\pm{0.02}$ & ---  \\\\ \\hline  \\end{tabular} \\end{center} \\end{table*}   \\begin{figure*} \\centering \\includegraphics[bb=14 14 562 484,width=15 cm,clip]{Figure2a.eps} \\caption{The H- and K-band continuum normalized ESO-NTT-SofI archival spectra of the stars WR20a (O3\\,If*/WN6+O3\\,If*/WN6), WR20b (WN6ha), WR21a (O3\\,If*/WN6+early O) together with that of the new stars WR20aa and WR20c. The main H, He and N emission lines are identified by labels.} \\label{FigVibStab} \\end{figure*}  Westerlund 2 (Wd2, \\citet{b27}, \\citet{b26}) is a young massive stellar cluster placed in the core of the H{\\sc ii} region RCW49, which is located in the fourth Galactic quadrant, projected in the area of the Carina constellation. This cluster hosts one of the most massive binaries known: the star WR20a, comprised by two O3\\,If*/WN6 type stars with absolute masses of 83M$_\\odot$ and 82M$_\\odot$ \\citep{b3,b19}. Another massive star of the same spectral-type, WR20b \\citep{b29} appears about 4' south-east of the cluster core. A second massive WR binary, WR21a (that is located 16' east of the Wd2 core) has a primary component of spectral type O3\\,If*/WN6 with an estimated minimum mass of 87M$_\\odot$, and an early O type secondary with a minimum mass of 53M$_\\odot$ \\citep{b1}. However the membership of WR21a to Wd2 is still uncertain. Such WR stars are members of the WNH type, a subset of very luminous and massive Hydrogen core burning WN objects (see definition in \\citet{b2} and references there-in).  \\begin{table*} \\caption{Journal of the spectroscopic data used in this work.} \\label{catalog} \\begin{tabular}{ccccccccc} \\\\ \\hline \\hline Star & Program ID &  Date     & Telescope & Instrument  &  Grism/Disperser   &   Slit  &  R  & Range ($\\mu$m) \\\\ \\hline WR20a & 075.D-0210 & 2005-06-21  & NTT       &    SofI     &  GR       & 0.6$\\arcsec$ x 290$\\arcsec$ & 1000 & 1.53-2.52  \\\\ & 072.D-0082 & 2004-01-29  & NTT       &    EMMI/2.3     &  GRAT\\#3 & 1$\\arcsec$ x 120$\\arcsec$ & 1600 & 0.39-0.47  \\\\ WR20b & 075.D-0210 & 2005-06-20  & NTT       &    SofI     &  GR       &  0.6$\\arcsec$ x 290$\\arcsec$ & 1000 & 1.53-2.52  \\\\ WR21a & 075.D-0210 & 2005-06-20  & NTT       &    SofI     &  GR       &  0.6$\\arcsec$ x 290$\\arcsec$ & 1000 & 1.53-2.52  \\\\ WR20aa & 075.D-0210 & 2005-06-21  & NTT       &    SofI     &  GR       &  0.6$\\arcsec$ x 290$\\arcsec$ & 1000 & 1.53-2.52  \\\\ & GOSSS & 2010-06-30  & du Pont      &    Boller \\& Chivens  &  1200 lines mm$^{-1}$       & 1$\\arcsec$x270$\\arcsec$ & 2500 & 0.39-0.55  \\\\ WR20c & 075.D-0210 & 2005-06-21  & NTT       &    SofI     &  GR       &  0.6$\\arcsec$ x 290$\\arcsec$ & 1000 & 1.53-2.52  \\\\ WR25 & GOSSS & 2008-05-20  & du Pont      &    Boller \\& Chivens  &  1200 lines mm$^{-1}$     & 1$\\arcsec$x270$\\arcsec$ & 2500 & 0.39-0.55  \\\\ HD93129A & GOSSS & 2010-06-30  & du Pont      &    Boller \\& Chivens  &  1200 lines mm$^{-1}$     & 1$\\arcsec$x270$\\arcsec$ & 2500 & 0.39-0.55  \\\\ \\hline \\hline \\end{tabular} \\end{table*}   With empirical masses exceeding 80 M$_\\odot$, WNH stars are probably among the most massive stars known, being  in the top-end of the mass distribution in young clusters and OB associations. This assumption is supported by results obtained from several studies of WNH stars in massive clusters such as R145 in 30 Doradus \\citep{b20}, WR25 \\citep{b22} and WR22 \\citep{b23} in the Carina Nebula region, and NGC3603-A1 in NGC3603 \\citep{b21}.  More recently, \\citet{b14} from the spectroscopic analyses of four WN stars located within R136 (the core of 30 Doradus Nebula in the Large Magellanic Cloud - LMC) and three WN stars of HD 97950 (the core of NGC 3603 in the Milky Way), computed initial masses in the range $165-320$ M$_\\odot$, and $105-170$ M$_\\odot$, for the stars in R 136 and NGC 3603, respectively. They point out that due to their proximity to the Eddington limit, the very high-mass progenitors would possess an emission-line spectrum, at the beginning of their main-sequence evolution, mimicking the spectral appearance of classical WR stars.  The spectral-type O3If*/WN6-A was introduced almost thirty years ago by \\citet{b28} to classify the bright emission line star Sk-67 22 in LMC, which shows an intermediate spectrum between those of HD93129A (O2\\,If*) and HD93162 (=WR25, O2\\,If*/WN6). New examples of this intermediate spectral type were mainly found in LMC, always consisting in bright massive objects (see for example, \\citet{b35}, \\citet{b36}, \\citet{b14}).    \\begin{figure*} \\vspace{0pt} \\centering \\includegraphics[bb=14 24 580 514,width=12 cm,clip]{Figure3b.eps} \\caption{The optical spectrum of WR20aa, together with that for HD 93129A (O2\\,If*), WR 25 (O2.5\\,If*/WN6 + O), and WR20a (O3\\,If*/WN6 + O3\\,If*/WN6). A comparison with the spectrum of HD93129A and that for WR20a shows that the optical spectrum of WR20aa approximates better to the later.} \\label{FigVibStab} \\end{figure*}   \\begin{table*} \\caption{Equivalent widths (\\AA\\ ) of the main emission lines detected in the NTT-SOFI H- and K-band spectra of WR20aa, WR20c WR20a, WR20b and WR21a.} \\label{catalog} \\centering \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{ccccccccccc} \\hline \\hline & Line  & He{\\sc ii}+Br12 &He{\\sc ii}+Br11 &He{\\sc ii}  &He{\\sc i}  & He{\\sc ii}+Br10 & He{\\sc i}+N{\\sc iii} & He{\\sc i}+He{\\sc ii}+Br$\\gamma$ &He{\\sc ii} \\\\ \\hline & $\\lambda$ ($\\micron$)  & 1.645  & 1.685 &  1.693   & 1.702 & 1.738 & 2.115 & 2.167 & 2.190 \\\\ \\hline & WR20b (WN6ha)  & -7.4 $\\pm$1.5 & -11.8 $\\pm$2.1 & -5.3 $\\pm$1.0 & -7.0 $\\pm$0.8 & -18.6 $\\pm$2.1 & -12.1 $\\pm$1.9  & -74.5 $\\pm$10.6  & -33.7 $\\pm$4.5 \\\\ & WR20a (O3\\,If*/WN6+O3\\,If*/WN6) & -3.0 $\\pm$0.7 & -3.3 $\\pm$0.7 & -1.7 $\\pm$0.4 & -1.3 $\\pm$0.3 & -8.0 $\\pm$1.4 & -2.0 $\\pm$0.5  & -29.9 $\\pm$6.3  & -14.5 $\\pm$3.8 \\\\ & WR21a (O3\\,If*/WN6+O) & -1.5 $\\pm$0.3 & -2.3 $\\pm$0.5 & -2.9 $\\pm$0.7 & -1.9 $\\pm$0.5 & -10.2 $\\pm$2.4 & -1.2 $\\pm$0.5  & -20.1 $\\pm$5.2  & -5.3 $\\pm$1.3 \\\\ & WR20aa (O2If*/WN6)& -3.1 $\\pm$1.0 & -3.3 $\\pm$1.0 & -0.2 $\\pm$0.1 & -1.6 $\\pm$0.4 & -6.1 $\\pm$1.3 & -1.1 $\\pm$0.4  & -9.9 $\\pm$2.1  & -4.8 $\\pm$1.2 \\\\ & WR20c (O2If*/WN6) & -1.5 $\\pm$0.5 & -1.5 $\\pm$0.6 & -1.3 $\\pm$0.5 & -1.3 $\\pm$0.4 & -5.2 $\\pm$1.2 & -2.7 $\\pm$0.7  & -12.5 $\\pm$2.7  & -4.0 $\\pm$1.2 \\\\ \\hline \\end{tabular} \\end{table*}  In this paper we report the identification of two new Galactic O2\\,If*/WN6 stars (WR20aa and WR20c) in the outskirt of the massive young cluster Wd2. In section 2 we describe the used observational data and the data reduction procedures, in section 3 we present the results and discussion, and in section 4 it is performed a summary of the work.    \\begin{table*} \\caption{Equivalent widths (\\AA\\ ) of the main diagnostic lines detected in the Blue optical spectra of WR20aa, HD93129A, WR20a, and WR25 stars.} \\label{catalog} \\centering \\renewcommand{\\footnoterule}{}  % \\begin{tabular}{cccccc} \\hline \\hline ID  & N{\\sc iv} 4058 &He{\\sc ii} 4686  &H$_\\beta$  & N{\\sc v} 4604 & N{\\sc iii} 4634-40-42 \\\\ \\hline WR25  & -2.5 $\\pm$0.1 & -15.6 $\\pm$0.7 & -2.7 $\\pm$0.1 & 0.3 $\\pm$0.1 & -0.4 $\\pm$0.1 \\\\ WR20a  & -2.3 $\\pm$0.1 & -14.0 $\\pm$0.6 & --- & 0.4 $\\pm$0.1 & -0.2 $\\pm$0.1  \\\\ WR20aa  & -2.4 $\\pm$0.1 & -9.6 $\\pm$0.5 & 0.3 $\\pm$0.1 & 0.7 $\\pm$0.1 & -0.2 $\\pm$0.1 \\\\ HD93129A & -1.3 $\\pm$0.1 & -3.4 $\\pm$0.4 & 1.3 $\\pm$0.1 & 0.5 $\\pm$0.1 & -0.1 $\\pm$0.1 \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\subsection{NIR spectra of WR20aa and WR20c}  In Figure 1 we present the telluric corrected (continuum normalized) H- and K-band spectra of the WR20aa and WR20c stars, together with those of WR20a, WR20b and WR21a. The strongest features are the blends of H{\\sc i} and He{\\sc ii} emission lines at 1.736$\\mu$m and 2.167$\\mu$m, as well as the He{\\sc i}+N{\\sc iii} and He{\\sc ii} emission lines at 2.115$\\mu$m and 2.189$\\mu$m, respectively. A list with the main emission lines and corresponding equivalent line-widths is presented in Table 3.  The similarity between the NIR spectra of the known O2\\,If*/WN6 stars with that of WR20aa and WR20c (Figure 1) is remarkable, indicating that the later may also belong to the O2\\,If*/WN6 spectral type. This assumption is reinforced by the close morphological match of the WR20aa H- and K-band spectra with that for WR20a, one of the most massive binaries known to date (O3\\,If*/WN6+O3\\,If*/WN6) for which masses of 83M$_\\odot$ and 82M$_\\odot$ were derived \\citep{b3,b32}, and by the match of the WR20c H- and K-band spectra with that of WR21a, another extremely massive binary system (O3\\,If*/WN6 + early O), for which \\citet{b1} estimated minimum masses of 87M$_\\odot$ and 53M$_\\odot$, respectively.  \\citet{b39} showed that some O4\\,If stars (like HD\\,16691 and HD\\,190429) can be erroneously classified as WN stars if observed only in the K-band. In fact, their K-band emission features are albeit morphologically identical to those seen in the spectra presented in Figure 1. However, as can be seen in \\citet{b40}, at least in the case of HD\\,190429 this is not true for the H-band. Indeed, the H-band spectrum of HD\\,190429 is completely different of those shown in Figure 1, being virtually featureless in the mentioned spectral range. Also \\citet{b38} present the H-band spectrum of Cyg OB2 \\#7 (O3\\,If*) which shows several H and He lines in \\textit{absorption}. On the other hand, as can be noticed from Figure 1 the H-band spectra of WR20aa and WR20c are \\textit{not} featureless, with all H and He relevant lines clearly seen in \\textit{emission}. In this sense, H- and K-band spectra of heavily reddened massive star candidates may became useful tools to discriminate between the OIf* and OIf*/WN spectral types.  \\subsection{The optical spectrum confirms: WR20aa is an O2\\,If*/WN6 star}   The spectral morphology derived from NIR spectrograms of one of the new sources (WR20aa), is confirmed from new optical spectroscopic observations. In Figure 2, we present the optical spectrum of WR20aa, together with GOSSS spectra of HD 93129A (O2 If*) and WR25 (O2.5\\,If*/WN6 + O) obtained with the same instrumentation and setup, along with the archival ESO/EMMI spectrum of WR20a (O3\\,If*/WN6 + O3\\,If*/WN6) \\citep{b46}. Table 4 lists the equivalent widths (EW) for the main diagnostic lines used for spectral classification of WR20aa, HD 93129A, WR20a, and WR25.  The blue spectrum of WR20aa is dominated by He{\\sc ii} 4686\\AA\\ and N{\\sc iv} 4058\\AA\\ emission lines, together with less pronounced emission lines of N{\\sc iii} 4634-40-42\\AA\\ and Si{\\sc iv} 4089\\AA\\ . Numerous absorption lines of H{\\sc i} and He{\\sc ii}, and also some diffuse interstellar bands (DIBs) at $\\lambda\\lambda$ 4428\\AA\\, 4726\\AA\\, 4765\\AA\\, and 4865\\AA\\ are present. A special mention deserve the strong N{\\sc v} 4604-20\\AA\\ absorptions, which are characteristic of the earliest O-type (O2-3) and WNha stars. These absorption lines show weak P-Cyg profile structure. The  strength of the N{\\sc iv} 4058\\AA\\ emission line relative to the N{\\sc iii} 4634-40-42\\AA\\ emission lines indicates that the spectral type of WR20aa is at least as early as O2, according to the criteria described in \\citet{b43}. In addition, the presence of a weak P-Cyg profile in the H-beta line of WR20aa, compared to the same line in HD\\,93129A, which shows no emission, and the P-Cyg profile observed in  WR25, is indicative of an intermediate spectral type between  OIf* and WNha. In this way, \\citet{b42} suggest that genuine WN6 stars are those with EW of He\\,{\\sc ii} 4686\\AA\\ lower than -12\\AA\\ . The measurement of He\\,{\\sc ii} 4686\\AA\\ line in WR20aa (-9.6\\AA\\, see Table 4) also indicates an intermediate spectral type. On the other hand, no detection of He\\,{\\sc i} absorptions is also indicative of a spectral type earlier than the O3. This picture is seen in the direct comparison with the spectra of HD\\,93129A and WR20a, which shows that the blue spectrum of WR20aa resembles better to the later (as it was already noticed from the NIR regime). It is also noticeable from Table 4 that the emission lines of WR20aa are closer in strength to those from WR20a, being the lines from the later slightly broader (this difference is perhaps related to the binary nature of WR20a). Consequently, we classify the spectrum of WR20aa as O2If*/WN6 \\citep{b46}.  A spectral classification for WR20c using optical MK criteria is not possible in this moment due to the lack of a good signal-to-noise optical spectrum of this star. From inspection of our NIR H- and K-band spectra, we find that WR20c is spectroscopically very similar to WR20aa, consequently we propose that WR20c should also be classified as O2\\,If*/WN6.   \\begin{figure*} \\vspace{2pt} \\centering \\includegraphics[bb=14 14 558 641,width=15 cm,clip]{Figure1.eps} \\caption{A mosaic made from the 3.6$\\mu$m Spitzer IRAC images (taken from the GLIMPSE survey archival data) of the region towards the RCW49 complex (6 arcmin $\\approx$ 14 parsecs at an heliocentric distance of 8 kpc). Here we identify the stars WR20a, WR20b and WR21a together with the two newly found O2\\,If*/WN6 objects, WR20aa and WR20c. We notice that a line connecting WR20aa and WR20c intersects the cluster centre just where the surface stellar density hits a maximum..} \\label{FigVibStab} \\end{figure*}  \\subsection{The new O2\\,If*/WN6 stars and the Westerlund 2 cluster: are they related?}  In Figure 3 it is shown a mosaic made from the 3.6$\\mu$m Spitzer IRAC images \\citep{b33} of the region towards the RCW49 complex. There we identify WR20a, WR20b and WR21a together with WR20aa and WR20c. It is interesting to notice that the three already known WR stars are well displaced from the core of the Wd2 cluster, which harbors at least a dozen of early O-type stars \\citep{b10}. The closest of them (WR20a), is located at about 0.6 arcmin from the cluster centre, with the others being placed at core distances of about 3.7 arcmin and 15.6 arcmin, respectively. The situation for WR20aa and WR20c is similar to that of WR21a, with the stars localized at angular distances (from the cluster centre) of 15.7 arcmin and 25 arcmin, well beyond the Wd2 cluster core. The presence of such massive stars well isolated in the outskirt of the RCW49 complex, naturally leave us to ask: how did they arrive there?  \\subsubsection{Interstellar extinction and absolute visual magnitudes }  The assumption of a possible physical connexion between WR20aa and WR20c with the Wd2 cluster population, in principle can be tested by comparing the (B-V) colors of these new stars with those of the presently accepted cluster members (WR20a and the 12 known early O-type members), which we do in Table 5. For the sake of completeness, we also included here the colors of WR20b and WR21a, the other two known WR stars in the vicinity of Wd2. The (B-V) colors of WR20aa and WR20c were computed from the B- and V-band photometry shown in Table 1, while the values for the other stars were taken from \\citet{b10} (WR20a and the early-O stars), \\citet{b30} (WR20b), and \\citet{b1} (WR21a). In the case of the early O-stars, we indicate the interval of (B-V) colors shown in Table 1 of \\citet{b10}. The corresponding visual extinctions were computed considering intrinsic colors (B-V)$_0$= -0.3 magnitudes, and the canonical value for the ratio of total to selective extinction R$_V$=3.1. Finally, absolute visual magnitudes were estimated assuming a distance of about 8 Kpc \\citep{b31,b10}.  In case of WR20aa we notice that its derived (B-V) color and visual extinction are compatible with those for the other massive stars in Wd2, while a very red color index (B-V)=2.6 was obtained for WR20c, which corresponds to about twice of the visual extinction inferred for the early-O stars of Wd2. The higher interstellar extinction derived for WR20c is probably due to the presence of a foreground molecular cloud as can be inferred from an inspection of the 2MASS J-band source density map\\footnote{http://aladin.u-strasbg.fr/java/nph-aladin.pl} of the region towards Wd2. Also, we remember that in the case of WR20c, the star is almost on the Galactic Plane which increases the probability that a dusty screen is on the line-of-sight of this star. In fact, it is not unusual that star forming regions present highly variable visual extinctions at scales of a few arcminutes. As examples, we can mention M17 \\citep{b44} and the Cygnus OB2 association \\citep{b45} in which the massive stars present a wide range (A$_V$\\,=\\,5-12 and A$_V$\\,=\\,5-20 magnitudes respectively) of individual visual absorptions.  On the other hand, the absolute visual magnitudes derived for WR20aa and WR20c are similar to that for WR20b, which is the only one in the area of RCW49, not known to be a binary. Indeed, if we correct the absolute visual magnitudes of WR20a (O3\\,If*/WN6+O3\\,If*/WN6) and WR21a (O3\\,If*/WN6+ early-O) in order to take into account the binary companion contribution (by 0.7 and 0.5 magnitudes, respectively), we can see that WR20aa is probably as luminous as the other OIf*/WN or WN6ha stars in Wd2, with WR20c being a bit less luminous than them. This is not surprising if one considers that stars of same specific spectral types may span a range in absolute magnitudes (see Table 3 of \\citet{b46}).  \\subsubsection{Two runaway stars?}  It is an observational fact that very massive stars preferentially are found in the core of their parental clusters, normally forming binary or multiple systems. However, there are a certain number of very young and massive stars that are found well isolated in the field. Two canonical scenarios try to explain the origin of these stars. The first known as the \\textit{binary-supernova scenario}, considers the disruption of a short-period binary system from the asymmetric supernova explosion of one of the binary components, while the second, named \\textit{dynamical ejection scenario}, involves dynamical three- or multiple-body encounters in dense stellar systems.  The time scale for the binary-supernova scenario involves values larger than the expected age (1-2 Myrs) for WN6ha stars \\citep{b14} although lower mass counterparts could reasonably be rather older (approx. 2.5 Myr). On the other hand, the multiple ejection of massive stars from dense massive clusters is not a new idea (see discussion in \\citet{b17}), and recent discoveries of very massive runaway stars (e.g. 30 Dor \\#16) highlighted the importance of this scenario in the evolution of massive clusters \\citep{b13,b18}. For more details on the two models see for example the discussion in \\citet{b11} and references there-in.  We searched for some observational constrain that could give support to the scenario in which WR20aa and WR20c were ejected from the Wd2 cluster centre following the dynamical ejection model. We found the surprising result that a line connecting WR20aa and WR20c (see Figure 3) intercepts the cluster core \\textit{exactly} where the star surface density hits a maximum. Also interesting, a vector conecting WR21a and the Wd2 core (see Figure 3) forms almost a right angle with the line connecting WR20aa and WR20c. This apparent configuration strongly resembles the dynamical scenario proposed by \\citet{b12} for the origin of AE Aur, $\\mu$ Col and the very eccentric binary $\\iota$ Ori, as runaway stars. From N-body simulations the authors conclude that these stars were ejected from their parental cluster as result of binary-binary interactions occurred in the Trapezium cluster. In this sense, it is remarkable to point out that the massive binary system WR21a is known to present a high eccentricity ($e\\sim0.64$, \\cite{b1}).  The fine geometrical alignment between WR20aa, WR20c and the Wd2 core, suggests that these stars possibly had a common origin somewhere in the central part of the cluster, being ejected from their birthplace (in a timescale not greater than their own age of no more than 2\\,Myrs \\citep{b14}, with minimum recession velocities (projected into the sky plane) of 18\\,km\\,s$^{-1}$ and 28\\,km\\,s$^{-1}$, respectively. It is also interesting that \\citet{b11} performed numerical simulations of a dynamical encounter between single massive stars and the very massive binary WR20a, founding that in a fly-by encounter the average recession velocity attained by a 70-80\\,M$_\\odot$ star could be quite moderate (less than 30\\,km\\,s$^{-1}$), being not formally classified as runaway.  Taking into account that the O2\\,If*/WN6 type stars are very rare and the perfect alignment of the WR20aa-cluster core-WR20c system, we postulate that both stars probably were formed somewhere in the core of the Wd2 cluster, being ejected from their birthplace by dynamical interaction with some other very massive star, in some previous stage of the cluster evolution. Of course, further observational (radial velocities, proper motions), and theoretical dynamical studies are still necessary to properly confirm our assumption. However, it is our opinion that the impressive alignment seen between such rare stars and the Wd2 cluster core, is in fact a strong clue favoring our hypothesis. In this sense, we speculate that WR20aa and WR20c \\textit{may be} two of the most interesting Galactic runways known to date."
    },
    "1107/1107.2598_arXiv.txt": {
        "abstract": "{The impact of quasars on their galaxy neighbors is an important factor in the understanding of the galaxy evolution models.} {The aim of this work is to characterize the close environments of quasars at low redshift (z$<$0.2) with the most statistically complete sample up to date using the seventh data release of the Sloan Digital Sky Survey.} {We have used 305 quasar-galaxy associations with spectroscopically measured redshifts within the projected distance range of 350 kpc, to calculate how surface densities of galaxies, colors, star-formation rates, oxygen abundances, dust extinction and ionization changes as a function of the distance to the quasars. We also identify and exclude the AGN from our main galaxy sample and calculate surface density of different galaxy types. We have done this in three different quasar-galaxy redshift difference ranges $|\\Delta$z$|<$ 0.001, 0.006, and 0.012.} {Our results suggest that there is a significant increase of the galaxy surface density around quasars, indicating quasar formation by merging scenario. We also observe a gap in the surface densities of AGN and blue galaxies at a distance of approximately 150 kpc from the quasar. We see no significant changes of star formation rate, color and ionization as function of distance from the quasar. From this investigation we cannot see any effects from quasars on their galaxy companions.} {} ",
        "introduction": "The extremely luminous nature of quasars is believed to be driven by an active galactic nucleus due to accretion of material upon a supermassive black hole \\citep[e.g.][]{Rees1984,LyndenBell1969}. Feeding such a black hole should affect the formation of galaxies in proximity to quasars \\citep{SR1998,Gunn1979}, and this could play a big role in the structure formation scenarios derived from the Standard Cosmology.  Could quasars be the results of mergers according to a hierarchical structure formation scenario predicted by Cold Dark Matter Models? For many years there have been controversial indications that both nuclear activity and star formation may be triggered or enhanced by interactions between galaxies. For instance, the presence of post-starburst populations in close quasar companions \\citep{Canalizo1997}, morphological asymmetries with high star formation rates near the quasar host \\citep{Gehren1984,Hutchings1984} and possible quasar companion galaxies \\citep{Canalizo2001}. Ionized gas is also observed around quasars \\citep[e.g.][]{Boroson1985}, and some authors have suggested that the quasars could even ionize the galactic medium up to several megaparsecs away \\citep{Babul1991,Rees1988}.  However, all these evidence-supported models and theories need much more substance if one wishes to raise their current status from speculation to fact. A great deal of work has gone into the study of close galaxy-galaxy pairs (less than 30 kpc) \\citep[e.g.][]{Rogers2009} and of the surrounding environments of quasars at larger radii (less than 1Mpc). Many studies of the quasar environment using the Sloan Digital Sky Survey (SDSS) database used galaxies with photometrically determined redshifts having large redshift uncertainties $\\delta$z $>$0.025. These allowed more quantitative studies than qualitative, which underlines the importance of careful quasar-galaxy association studies with spectroscopically defined redshifts $\\delta$z $<$ 0.001.  I have used a non-volume-limited sample drawn from the seventh data release (DR7) of the SDSS, containing more than 100,000 quasars, to study how quasars influence their nearby galaxies at low redshifts. In my thesis I examine how colors, surface densities, star formation rates, ionization of the galactic medium and certain spectral lines of galaxies near quasars change depending on the distance between the quasar and the neighbor. I have also investigated the active galactic nuclei (AGN) activity in host galaxies near quasars. This study is done for a sample consisting of 305 quasar-galaxy associations at redshifts 0.03 $<$ z $<$ 0.2 using Standard Cosmology ( $\\Omega_{\\Lambda}$=0.70, $\\Omega_{M}$=0.30, H$_{o}$=70 km/s/Mpc.) I hope to see in what ways quasars may impact surrounding galaxies, which is of major importance to understand AGN-feedback models and the formation of the most luminous structures in our Universe. ",
        "conclusions": "We have used the Sloan Digital Sky Survey (SDSS) and utilized spectral lines, dereddened apparent magnitudes, redshifts and angular distances of 305 quasar-galaxy associations at redshifts 0.03 $<$ z $<$ 0.2 and projected distances less than 350 kpc, to study quasar neighborhoods.  We corrected the colors and spectral lines for dust extinction in the non-AGN companion galaxies. We calculated colors, surface densities, ionization ratios F([\\ion{O}{iii}]/[\\ion{O}{ii}]) and F([\\ion{O}{iii}]/H$\\beta$), oxygen abundances and star-formation rates (SFR) from H$\\alpha$ and [\\ion{O}{ii}] fluxes for the non-AGN galaxies. We also estimated the surface density distribution of AGNs in our galaxy sample using BPT-diagram combined with Kauffmann's criteria \\citep{Kauffmann2003} and the H$\\alpha$ width. To see how different properties of the companion galaxy vary as function of distance from the quasar, we binned the surface and number densities in bins of 35 kpc and the rest of the above parameters in bins of 50 kpc, the latter in order to be easier comparable to similar studies with galaxy-galaxy pairs \\citep[e.g.][]{Nikolic}. We excluded the AGN in our non-AGN-galaxy neighbor sample, and calculated the average parameter value and error bar for all galaxies within each bin. We plotted the average values of color, SFR, ionization parameters, dust extinction and oxygen abundance in each bin as function of distance between galaxy and quasar. Finally, we plotted LSAs of the same parameters as function of distance from the quasars, and estimated the maximum correlation from three different weights. Some conclusions are:  \\begin{enumerate}  \\item No correlation between color and distance in the binned plots, while the LSAs give a maximum 1$\\sigma$ decrease of $U_{e}-R_{e}$ towards the bluer near the quasars.  \\item We cannot see any difference in dust extinction depending on distance from quasar in the bins. Partially, this might be a problem of having few galaxies with well-measured H$\\beta$, which stresses the importance of further studies. However, the LSAs show a maximum 2$\\sigma$ decrease of dust near the quasars.  \\item There is a non-agreement between the H$\\alpha$- and [\\ion{O}{ii}]-derived star-formation rates in both the binned plots and the LSAs. In the LSAs, no change of H$\\alpha$-SFR is seen at decreased distance from the quasars, while the [\\ion{O}{ii}]-derived SFRs show an increase 1$\\sigma$ $<$ trend size $<$ 2$\\sigma$ near the quasars.  \\item No difference in ionization is seen whatsoever, neither by investigating F([\\ion{O}{iii}]/[\\ion{O}{ii}]) nor F([\\ion{O}{iii}]/H$\\beta$) in the bins, while LSAs can show a 0 $<$ trend size $<$ 4$\\sigma$ increase near quasars.  \\item There is a too large scatter in the oxygen abundance plot in order to see any correlation of the oxygen abundance with distance. A max 1$\\sigma$ decrease can be seen in the LSAs.  \\item There is a gap at the distance of 150 kpc in the number of AGN and as well as in the number of blue galaxies, something that might depend on too low number to make good statistics from, or more interesting physics. Further studies with increased number of galaxies are of very high importance.  \\item The most important conclusion we have is a large increase in the surface density of all galaxies closer than 100 kpc to the quasars in both samples, which could support a merger scenario behind quasar activity. In case mergers are responsible for quasar activity, perhaps could some of the AGNs in the galaxy sample can be secondary products from the same mergers? \\end{enumerate}  \\subsection{Greatly enhanced surface densities}  We investigated the number densities of non-AGN companion galaxies. In the closest bins in column 1 we see a significant increase in the number of galaxies, especially blue galaxies. From the same figure we can also see that when using larger $|\\Delta z|$ cuts (0.006 and 0.012) the relative number of galaxies in the furthest bins increase significantly. This is a result of including more background and foreground galaxies, which has a considerably higher probability for greater projected distances \\citep{Nikolic} at the larger $|\\Delta z|$ cuts. We had more red galaxies than blue galaxies in our non-AGN-neighbor sample, especially further out in projected distance.  Histogram binning is sensitive to the number of bins chosen. Using a wrong number of bins can make ordinary noise appear as patterns in histograms and giving space for misinterpretations in the analysis process. To avoid this we tried using a histogram binning program that calculates the optimal bin size \\citep{Knuth}, but since we have too few galaxies it turned out to be not applicable for this problem.  In a previous chapter we plotted the surface densities of the blue neighbor galaxies, the red galaxies and all galaxies combined and found a huge increase in surface density of all galaxy types at bins closer than 150 kpc to the quasar. Especially high is the increase in blue gas-rich galaxies and AGN if one compares to the background of galaxies further away from the quasar (d $>$ 150 kpc). In case mergers are responsible for quasar activity, perhaps the AGN and blue galaxies could be secondary products from the same mergers?  Other authors \\citep[e.g.][]{Strand2008, Serber2006} also found an increase at close projected separations in the number of galaxies around quasars at low redshift. However, the large scales used in those studies (up to 2 Mpc) give a large contamination of non-physical pairs at projected separations larger than 150 kpc \\citep{Nikolic} that increases steadily at greater separations, and with larger $|\\Delta z|$ between the objects in the pair \\citep{Strand2008}.  At the same time an increased density around the quasars does not necessarily imply a merger scenario, even though the temptation to assign this explanation to the result might be overwhelming. Monolithic collapse models can predict the formation of smaller objects around more massive ones when both smaller and larger objects collapse in the gravitational well of the same halo \\citep[e.g.][]{Press}. The LCDM model itself predicts a large number of satellite galaxies around massive galaxies, something that is still difficult to know whether it is the case or not.  How could we test whether our surface densities are products of hierarchical structure formation or of monolithic collapse? One way would be to construct a sample of field quasars at low-redshifts (0.03 $<$ z $<$ 0.2) but that have no companion the closest 350 kpc. In case of merger scenario being the secret of quasar activity, these quasars should have an average luminosity that is fainter than in the sample with quasars that have companion galaxies. An second way is by using modelling to predict the number of lower-mass objects around quasars of the average luminosity in our quasar sample \\citep[e.g.][]{Carlberg1990}. Similar attempts of investigating the clustering of quasars have been done \\citep[e.g.][]{HopkinsCluster} where an excess of clustering could be shown, but where the results could not be clearly put in context of any given structure formation scenario. A third way, is by using spectral synthesis modeling for tracing the star formation histories of the companion galaxies. This could reveal signs of previous mergers.  \\subsection{Comparison to galaxy-galaxy pairs}  One way of identifying if our observed phenomenons are related to the quasars themselves or are normally present for all kind of galaxies, is by comparing to the environment of field galaxies.  We also generated surface and number density plots of field galaxy environments with similar absolute magnitude distribution to our quasars, in order to directly compare the environment of galaxies with same luminosity as our quasars and in the same kind of environment as the quasars namely galaxies in the field. Unfortunately, our sample size of field galaxy pairs was far too small to make any detailed or firm conclusions, and only the number and surface densities of the $|\\Delta$z $|<$ 0.012 that are independent of the colors of the satellite galaxies can be used for comparison. From these plots we can see that there seem to be a small, but no great increase of satellite galaxies around the main field galaxies, indicating that they were probably not formed in the same ways as the quasars. A similar increase at close distances in galaxy-galaxy number density would have supported that any galactic objects have tendencies to merge sooner or later or that the pair of galaxies got simultanously created in the same halo. Our sample size is although far too small to be reliable.  \\begin{figure*}[htb!] \\centering \\includegraphics[width=16cm,height=10.5cm]{ggNDtot.eps} \\caption{Number distribution of galaxies around the central field galaxies. Left panel shows blue galaxies. Middle panel shows red galaxies. Right panel shows all non-AGN galaxies. Three different redshift difference cuts are pictured for the pairs.} \\label{ggNDtot} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=16cm,height=10.5cm]{ggNFig85.eps} \\caption{Annular surface densities of different types of galaxies around central field galaxies. Surface density distribution of galaxies. In the a) left panel, blue galaxies with $U_{e}-R_{e}$$<$2.2, in b) first middle, red galaxies, $U_{e}-R_{e}$$>$2.2 and in c) the right middle panel, all non-AGN galaxies d) AGN. Three different redshift difference cuts are pictured for the pairs.} \\label{ggNFig85}% \\end{figure*}  A study with galaxy-galaxy pairs in the SDSS \\citep{Ellison} showed that in low-density environments the number density distributions are different from the number distributions of our quasar companions. This group used 5784 close galaxy pairs from the SDSS data release 4 (DR4) in different density environments with $zconf$ $>$ 0.7, projected distances within 80 kpc, redshift cut $|\\Delta$z $|<$ 0.017 and a redshift $<$ 0.1. Their number density distribution behaved differently from our surface densities from the quasar-galaxy pairs. In fact the numbers of companion galaxies in Ellison et al. were dropping off at closer distances, contrary to our case with non-AGN galaxy companions to the quasars. If only comparing their bins at 35 kpc and 70 kpc with ours, one can see that while they had a remarkable difference with a decrease in the closest bin (N$_{35}$/N$_{70}$=2/3), we have no difference in the number of galaxies in these two bins. The general increase in number density around quasars seems to be specific for quasars, which might be explained by the fact that quasars need to surround themselves with fuel to survive.  Maybe it is not so surprising that the environments of quasars and field galaxies look different. Coldwell \\& Lambas did study the surface density distributions of galaxies close to clusters, close to field galaxies and close to quasars. They could in their paper show that galaxies in the neighborhood of field galaxies have a lower density environment than galaxies around quasars do have.  \\subsection{Missing AGN and blue galaxies?}  Another curious thing about the surface density plots is the sudden drop of blue galaxies and AGN in the bin around 150 kpc. This could be a product of noisy data or low number statistics, but at the same time we see a slight increase of red galaxies in the 35 kpc closer surface density bin. Could this mean that those missing blue galaxies got trapped by the gravitational field of the quasar and got their gas stripped off (hence became red) as they fell into circular orbits around the quasar?  We remind the reader that the surface densities are shown on a logarithmic scale, and that it actually rather seems like the loss of blue galaxies in the 100-150 kpc bin is larger than the gain of red galaxies. Maybe some galaxies lost the gas when falling into circular orbits and got too faint to be detectable by the SDSS. Perhaps there is an unseen population of very faint galaxies in this bin. What if the galaxies simply not had a chance even to form any stars in their recent star formation history, before getting trapped by the quasars? A study of Milky Way satellites \\citep{Grcevich} has revealed that within 270 kpc of the Milky Way, almost all satellite galaxies have lost their gas and are spheroidal. This is strong support for tidal stripping processes within the distance range in our own study. What if something similar has happened to our blue non-AGN companion galaxies? Highly speculative scenarios would however need careful measurements of the velocity fields of the close quasar companion galaxies in order to follow how the dynamics of the gaseous content of the companion galaxies changes around the gap.  Many thoughts remain as pure speculations until we find a way to increase our neighbor galaxy sample without losing precision in redshift estimates. It is difficult to know whether the gap is real or simply an artefact from having too few quasar-galaxy pairs.  \\subsection{Influence on environment?}  There is an overall lack of clear correlations in color, F([\\ion{O}{iii}]/[\\ion{O}{ii}]), accretion rate F([\\ion{O}{iii}]), SFR, dust and F([\\ion{O}{iii}]/H$\\beta$). Overall, it looks like there might be a small increase in ionization and SFR at short projected separations as well as a very small decrease of color and dust. What makes this to be a less likely conclusion is the two deviating SFRs from H$\\alpha$ and [\\ion{O}{ii}]. We tried obtaining GALEX-data of FIR emission to get a third probe for star-formation in our galaxies, but FIR-emission data only existed for 27 of our neighbor galaxies, which was not usable in our study. A deviation as the one demonstrated for SFRs from H$\\alpha$ and [\\ion{O}{ii}] is likely to happen either if the dust extinction and/or the oxygen abundance changes with the distance to the quasars $or$ if we have more hidden AGN who could not be properly excluded at short projected separations. Since we have an overall increase of the AGN surface density, it is also likely that the number of hidden AGN is higher at close separations. This could lead to all the small trends observed: less dust (AGN can blow it away), no direct increase in SFR from H$\\alpha$, while a deviation from the SFR from the [\\ion{O}{ii}], and increase in ionization and a slightly bluer color. As we saw no dramatical changes of the dust extinction or oxygen abundance with the distance to the quasars, the second option with contamination of AGN seems more likely.  To try out the hypothesis, we tested by contaminating our sample with AGN and perform LSAs. Doing this can generates a 2$\\sigma$ decrease of $U_{e}-R_{e}$, up to 5$\\sigma$ increase of SFR from H$\\alpha$ (while still disagreeing with the SFR from [\\ion{O}{ii}]. This is especially the case when analysing with the weighting w$\\sim$ 1/($\\sigma)^{2}$, which means correlation analyses must be carried out with carefulness of both weighting and a proper AGN exclusion.  One could argue that any correlations might be buried in the redshift-dependence of e.g. SFR. We tried this by doing LSAs with different weighting models for SFR, ionization, color and dust extinction and divide it with the surface area, but could still not find any significant correlations.  The general lack of trends could also be a result from having a sample consisting of faint quasars. At the same time we would need to do a similar study with higher redshift quasars with galaxy companions to know whether a possible influence could be luminosity-dependent. Everything can however turn out to be different when one does a study with only very strong high-redshift quasars, and maybe only is AGN feedback important for these objects. Since the present day cataloges of quasars have too few observable companion galaxies we cannot do a similar high-redshift study now, but simply have to wait for the future. At the moment the SDSS DR7 does not allow either study due to the far too small sample size one obtains from the present sample and the lack of objects with spectroscopic redshifts above 0.2.  To predict whether any change in the degree of ionization actually should be observed with our faint quasars, one could estimate the radiation field from the quasars to see how the flux of high-energy photons changes at distances up to 350 kpc. We tried the alternative way of constructing a field galaxy sample of the same luminosity range to see how the satellites to these get influenced as a way to compare with, but the rather small sample made it difficult to state any conclusive arguments.  Another thing to take into consideration might be that the corrections for internal extinction we have used might be insufficient to show us any trends. We tried to perform the extinction corrections as accurate as possible by correcting the colors for extinction dependent of angle of view of the galaxy and by adjusting emission lines for extinction according to Whitford \\citep{Whitford}. Something else that might affect the results, is that a correlation might be buried in the redshift evolution of the parameters. We tried to take that into consideration by using the $|\\Delta z|$ cuts, with the lowest corresponding to the minimum spectroscopic redshift error.  The most important problem in the study was however still the small sample we had with many big individual errors in the fluxes we used, which caused many problems when investigating how different properties could eventually vary with the projected distance from the quasars. We tried surpassing the problem by increasing our sample with the 2MASS catalogue, but only three more quasar-galaxy associations could be found there, which was not helpful for the problem. We can only hope that the future will bring greater surveys with many more objects with careful line measurements, and therefore wait and hope for the DR8 and DR9 to fulfill our expectations."
    },
    "1107/1107.0762_arXiv.txt": {
        "abstract": "We explore the physics of shock evolution and particle acceleration in non-relativistic collisionless shocks using multidimensional hybrid (kinetic ions/fluid electrons) simulations. We analyze a wide range of physical parameters relevant to the acceleration of cosmic rays (CRs) in astrophysical non-relativistic shock scenarios, such as in supernova remnant (SNR) shocks. We explore the evolution of the shock structure and particle acceleration efficiency as a function of Alfv\\'enic Mach number and magnetic field inclination angle $\\theta$. We show that there are fundamental differences between high and low Mach number shocks in terms of the electromagnetic turbulence generated in the pre-shock zone and downstream; dominant modes are resonant with the streaming CRs in the low Mach number regime, while both resonant and non-resonant modes are present for high Mach numbers.  Energetic power law tails for ions in the downstream plasma can account for up to $15\\%$ of the incoming upstream flow energy, distributed over $\\sim5\\%$ of the particles in a power law with slope $-2\\pm0.2$ in energy. Quasi-parallel shocks with $\\theta\\le45^\\circ$ are good ion accelerators, while power-laws are greatly suppressed for quasi-perpendicular shocks, $\\theta>45^\\circ$. The efficiency of conversion of flow energy into the energy of accelerated particles peaks at $\\theta=15^\\circ$ to $30^\\circ$ and $M_A=6$, and decreases for higher Mach numbers, down to $\\sim2\\%$ for $M_A=31$. Accelerated particles are produced by Diffusive Shock Acceleration (DSA) and by Shock Drift Acceleration (SDA) mechanisms, with the SDA contribution to the overall energy gain increasing with magnetic inclination. We also present a direct comparison between hybrid and fully kinetic particle-in-cell results at early times; the agreement between the two models justifies the use of hybrid simulations for longer-term shock evolution. In SNR shocks, particle acceleration will be significant for low Mach number quasi-parallel flows ($M_A < 30$, $\\theta< 45$). This finding underscores the need for effective magnetic amplification mechanism in SNR shocks. ",
        "introduction": "\\indent  Non-relativistic astrophysical shocks have long been considered as probable sources of acceleration of galactic cosmic rays (CRs) \\citep{Chen:1975p1774,Bell:1978p884,Blandford:1978p891}. The CR spectrum extends over many decades in energy, following a power law $f(E)\\propto E^{-\\alpha}$ distribution. Nonthermal emission from supernova remnant shocks and from relativistic jets in Gamma-ray bursts and Active Galactic Nuclei is also modeled as synchrotron or inverse Compton scattering from a power law population of electrons at these sources. While the shape of the power law spectrum can be explained by modern shock acceleration theory, particle acceleration efficiency, i.e., the number and energy fraction contained in the accelerated particles, depends on nonlinear shock physics and has to be constrained through kinetic  simulations. Such simulations are the subject of this paper.  The first-order Fermi acceleration mechanism at shocks \\citep{Bell:1978p884,Blandford:1978p891,Blandford:1987p1773} is the strongest contender to explain the observed power law spectra. The mechanism works by having particles cross a shock front several times, with a positive energy gain on every crossing. Diffusive Shock Acceleration (DSA) theory describes how these particles scatter upstream and downstream from the shock due to electromagnetic (EM) turbulence, and predicts a universal power law slope $\\alpha$, which to zeroth order depends only on the shock compression ratio ($n_d/n_u$, the ratio of the downstream to upstream density). Both the particle acceleration efficiency and the acceleration rate depend on the properties of turbulent fields and on the momentum of the accelerated particle; this is accounted for in the diffusion coefficient in the DSA model. Since the diffusion coefficient is not directly measurable in shocks of astrophysical origin, simplifying assumptions are usually made when estimating the maximum particle energies and acceleration times for a given object. One common assumption is to consider Bohm diffusion, by setting the diffusion coefficient to be proportional to the Larmor radius of the accelerated particle. This assumption then constrains the energy gain to be directly proportional to time, $E(t)\\propto t$, and also limits the maximum energy of a particle by setting the maximum Larmor radius $r_{Li}$ to be of the same order of magnitude as the system size.  Another relevant acceleration mechanism is Shock Drift Acceleration \\citep[SDA,][]{Chen:1975p1774,Webb:1983p1775,Begelman:1990p1777} and Shock Surfing Acceleration \\citep[SSA,][]{Lee:1996p3421,Hoshino:2002p3454,Shapiro:2003p3458}, where a particle is accelerated by the $\\vec{E}_0=-\\vec{V}\\times\\vec{B}$ upstream electric field in a magnetized shock, with $\\vec{B}$ the upstream magnetic field, and $\\vec{V}$ the upstream plasma flow velocity measured in the downstream frame. The main difference between SDA and SSA is the importance of the electrostatic cross-shock potential for the confinement of particles in the acceleration region; nevertheless, particles are accelerated by the same upstream electric field $E_0$, and for the purposes of this work we do not distinguish between the two. After being reflected at the shock, a particle will drift along the shock front in the direction of this upstream electric field, gaining energy. For magnetized shocks, this mechanism works along with DSA to produce energetic particles.  Astrophysical collionless shocks can be characterized by their Alfv\\'enic Mach number, defined as $M_A=V_{sh}/V_{A}$, where $V_{sh}$ is the shock front velocity (in the upstream frame) and $V_A=B_0/\\sqrt{4 \\pi n_0 m}$ is the Alfv\\'{e}n velocity (with $B_0$, $n_0$ and $m$ being the upstream magnetic field, ion number density and  ion mass, respectively). The magnetic field inclination angle $\\theta$ between $\\vec{B}$ and the shock normal also plays an important role. Generation of electromagnetic turbulence in the pre-shock zone is crucial for the DSA mechanism, since particles need to be scattered back to the shock front multiple times for efficient acceleration. Since cold particles will be constrained to move along the magnetic field lines, different shock configurations will develop for angles $\\theta\\le 45^\\circ$ (quasi-parallel shocks), and for angles $\\theta>45^\\circ$ (quasi-perpendicular shocks). The main difference for the two cases is that for $\\theta>45^\\circ$ there is a sharp decrease in the number of particles escaping from the shock. For quasi-parallel shocks an extended pre-shock region will develop, with waves and turbulence being generated by reflected ions that stream in front of the shock; quasi-perpendicular shocks will exhibit very little self-generated turbulence upstream, and instead a shock-foot region with the size comparable to the ion Larmor radius $r_{Li}$ will form. The DSA mechanism will then play a crucial role in particle acceleration in quasi-parallel shocks \\citep[for a comprehensive review of microphysics of collisionless shocks see][]{Treumann:2009p2090}.  Several approaches are commonly used to model the intrinsically nonlinear processes that dominate shock acceleration. Semi-analytic kinetic models \\citep{Kirk:1989p1778,Ballard:1991p1780,Achterberg:2001p1795} provide a theoretical background for the description of the particle acceleration problem. Similarly, Monte Carlo test particle simulations with prescribed turbulence models have been performed by several authors \\citep[e.g.,][]{Bednarz:1998p1798,Ellison:2004p1800,Niemiec:2004p1799}. These can be extended over very large periods of time and thus have predictive behavior over the long term evolution of particle acceleration. They lack the self-consistent properties of first principle simulations, such as fully kinetic particle-in-cell (PIC) models. In PIC models, the full set of Maxwell's equations is solved on a finite grid, where particles are also pushed and serve as source terms for the creation of EM fields \\citep{birdsall}. Hybrid simulations follow a similar method and consider kinetic ions and fluid electrons \\citep{lipatov}; since electron time-scales do not need to be resolved, evolution of the physical system for longer times is possible.  Previous work on multidimensional PIC shock simulations has been conducted for relativistic pair plasmas \\citep{Spitkovsky:2005p1892,Spitkovsky:2008p1909,Sironi:2009p1330,Nishikawa:2009p3515}, and for relativistic ion-electron plasmas \\citep{Spitkovsky:2008p1910,Martins:2009p1933,Sironi:2011p3059}. In the non-relativistic regime, both PIC \\citep{Amano:2007p2166,Kato:2008p2139,Amano:2009p2160,Amano:2010p2162,Riquelme:2011p3505} and hybrid simulations exist \\citep{Bennett:1995p1175,ELLISON:1993p40,GIACALONE:1992p45,GIACALONE:1994p41,GIACALONE:1997p1192,SCHOLER:2002p2491,Giacalone:2004p44,Sugiyama:2011p3494}, among others. Electron acceleration is found to be more efficient in (magnetized) quasi-perpendicular shocks, in the non-relativistic case, while ion acceleration is mainly observed in quasi-parallel shocks.  In this work, we concentrate on the systematic analysis of non-relativistic magnetized ion-electron shocks for diverse Mach numbers and magnetic field inclination angles. We extend previous work by analyzing a wide range of parameters, relevant for both solar wind shocks and for SNR shocks, and follow the shock evolution and ion acceleration for very long times and extended regions of space in two dimensions. Such extensive surveys are not commonly found in the literature (\\citet{GIACALONE:1997p1192} presents a similar survey, done in 1D, and for a more limited range of Mach numbers). Qualitative differences are found between high and low Mach number shocks, and for quasi-parallel and quasi-perpendicular shocks. The analysis of individual ion trajectories provides a direct measurement of how particles are accelerated, and allows us to discriminate between the DSA and SDA models, and their relative contributions to the high-energy end of the particle distribution functions. We focus on particle acceleration and how the main shock parameters affect the acceleration efficiency. We also analyze the particle spectra behavior in time, and compare between spectra for the different Mach numbers and B field inclination angles, providing a complete survey of ion acceleration in non-relativistic magnetized shocks.  This work is organized as follows. In Section \\ref{section2} we present the simulation setup and the physical parameters to be explored. In Section \\ref{section3} we show the global shock properties, including the main features of quasi-parallel and quasi-perpendicular shocks, together with typical particle spectra found in each case. In Section \\ref{section4} we present the results of the complete shock parameter survey, analyze the shock properties with particular emphasis on wave and turbulence generation, and its influence on the measured particle spectra. In Section \\ref{section5} we consider the particle acceleration mechanisms and discuss how their relative contributions vary with the shock parameters. The final discussion, consequences and constraints for astrophysical shock scenarios derived from the present work are presented in Section \\ref{section6}. Finally, we add an Appendix presenting a direct comparison between fully kinetic PIC simulations and their hybrid counterparts, and we include a detailed analysis of the numerical parameters used. ",
        "conclusions": "\\label{section6}  In this paper we have discussed ion acceleration in non-relativistic magnetized shocks. In an effort to understand particle acceleration, we conducted a complete 2D survey of this class of shocks, varying both the Alfv\\'enic Mach number and the magnetic field inclination angle for each shock. As shown before by other authors, our results confirm that there are fundamental differences between quasi-parallel and quasi-perpendicular shocks. Our study allows to systematically trace these differences as functions of the flow parameters and see the transition between different regimes. For quasi-parallel shocks, a significant number of particles are reflected by the shock back into the upstream medium, which causes electromagnetic instabilities that propagate back to the shock, and create a turbulent upstream medium. In contrast, for quasi-perpendicular shocks, with $\\theta \\ge 60^\\circ$, there is no significant particle population being reflected at the shock front, and hence EM turbulence upstream is greatly reduced.  Our results show a fundamental difference between shocks at low $M_A \\le 10$ Mach numbers, and shocks with $M_A > 10$; for low Mach number shocks a right-hand mode with $\\delta B/B_0<1$ is excited in the upstream, resonant with the ions streaming away from the shock, whereas for higher Mach numbers a left-hand non-resonant mode predominates. In the high Mach number case, the left-hand non-resonant waves exhibit B field amplitudes $\\delta B/B_0>1$; these waves resemble the CRCD instability, as they have the same polarization. Here however, the wavelength $\\lambda$ is comparable to the Larmor radius of the lowest-energy CR particles, and the phase velocity of the waves is around the Alfv\\'en velocity. Thus, from the point of view of the highest-energy CR particles, since $\\delta B/B_0>1$, $\\lambda<r_{Li}$ and $V_A \\ll V_{flow}$ (i.e.,the phase velocity of the waves is much less than the upstream flow velocity), the ion-whistler waves observed should have similar scattering properties to CRCD-generated electromagnetic turbulence.  For high Mach number shocks the self generated EM waves with $\\delta B/B_0>1$ in the upstream medium start to reflect particles locally, and the shock is observed to reform further upstream; a similar effect of shock reformation due to Short Large-Amplitude Magnetic Structures (SLAMS) has also been described in the literature \\citep{Dubouloz:1995p1593}. This shock jump is observed at a relatively early stage $t\\sim100\\,\\mathrm{\\omega_{ci}^{-1}}$ of the shock propagation, and it is thus possible that this is a transient effect due to the shock formation process. There are indications, from longer simulations (in time) over larger spatial domains that this is the case: although reformation still happens at late times, its effects are less noticeable than at the time the shock is forming. These simulations also show that the dominant EM modes change in time if the particle flux changes as well, which leads to a highly dynamic shock structure where sub-shocks form in localized regions of the upstream medium. Thorough analysis of such effects will be the subject of a future publication.  We compared shock acceleration efficiencies in similar conditions for all Mach numbers and angles, as described in Table \\ref{table1}, and showed that the highest energy-conversion efficiencies are achieved for intermediate angles $\\theta\\sim15^\\circ$ to $30^\\circ$ and low Mach numbers. The power law indexes are similar in most cases, yielding $\\alpha=2\\pm0.2$. The typical growth of maximum energy in time is found to be $E\\propto t^{0.28}$ (for $M_A=3.1$ and $\\theta=0^\\circ$). This value is indicative of a diffusion process less effective than Bohm diffusion, which would correspond to $E\\propto t$; a full comparison of diffusion coefficients for different Mach numbers and inclination angles will also be the subject of future work.  The analysis of representative particle tracks for two quasi-parallel low Mach number shocks is consistent with DSA mechanism, and confirms that particles are reflected multiple times towards the shock front. Overall, we can conclude that ion acceleration in magnetized non-relativistic shocks is efficient at finite intermediate angles, in the quasi-parallel regime, and for low Mach numbers ($M_A<10$ ). This is explainable if we take into account that injection into a diffusive acceleration process is more likely at low angles, since a significant number of particles are allowed to travel farther into the upstream medium, and generate EM waves and turbulence which enhance further particle reflection towards the shock. For higher inclination angles, however, a particle with a sufficiently high energy will be more easily confined to the shock front, and will be able to accelerate faster due to stronger shock drifting (Figure \\ref{fig9}); in terms of energy-conversion efficiency an intermediate angle yields the best results. Also, although higher Mach number shocks are able to accelerate particles to higher energies (as a fraction of incoming beam energy), the shock efficiency is lower because injection into the acceleration mechanism is less efficient at higher Mach numbers.  It is interesting to compare the present results with a similar shock survey done in \\citet{GIACALONE:1997p1192} (henceforth G97) using 1D hybrid simulations for typical solar wind flow parameters. For a more limited range of Mach numbers, from $M_A=1.5$ up to $M_A=12$ the authors also show downstream spectra of accelerated particles; the power-law tails extend over a similar range in energy as in our case, although high-energy particle splitting, and the introduction of external EM turbulence was used to facilitate the diffusive process in the case of G97. In the comparable range of Mach numbers (up to $M_A=12$), the conclusions about shock acceleration efficiency are similar. At higher Mach numbers (not explored in G97) we observe a clear decrease in efficiency, and different features such as shock reformation and the left-hand wave polarization mode.  Energetic arguments, in order to explain CR energies of up to $10^{15}\\,\\mathrm{eV}$ from SNR sources, usually require a large fraction $> 30 - 40\\%$ of the total energy in the shock to be transferred to accelerated particles, which indicates that shocks with $M_A < 30$ and $\\theta<30^\\circ$ are the most likely candidates for CR acceleration. For angles greater than $45^\\circ$, injection is very limited and no significant acceleration is observed; it should be noted, however, that if pre-accelerated particles could be injected in shocks with $\\theta\\sim60^\\circ$ in sufficient numbers, it would be possible to have higher acceleration efficiencies due to enhanced shock drift acceleration in the shock front (although this does not seem to occur naturally in the simulations).  Relatively low Mach numbers ($M_A < 30$) are thus required for effective ion acceleration, which then implies that strong magnetic field amplification should occur at SNR shock sites (if we consider typical outflow velocities of $\\sim1000\\,\\mathrm{km/s}$ in interstellar fields of $1\\,\\mathrm{\\mu G}$). This strong B field amplification is inferred from observations \\citep[for a review see][]{Reynolds:2011p3516}, and it might be achievable by the CRCD instability \\citep{Bell:2004p24,reville,Stroman:2009p2260,Riquelme:2009p2089,Gargate:2010p1604}. Such low Mach numbers are also favored by the electron acceleration mechanisms in shocks \\citep{Riquelme:2011p3505}. We do see strong B field amplification in our current simulations, but the effective Mach number near the shock front is at most reduced by half (effective Mach number near the shock is $M_A\\sim15$, corresponding $\\delta B\\sim2 B_0$, for the case of $M_A=31$, $\\theta=0$ far upstream conditions). A magnetic field amplification of at least an order of magnitude would be necessary in a SNR shock. As such, a self-consistent shock simulation where both accelerated particles and the CRCD instability are observed and where $\\delta B\\sim10 B_0$ is still lacking.  The numerical properties of the hybrid and PIC codes must be very well understood in order to validate results. We show the effects of dimensionality, and validate our choice of numerical parameters in the Appendix. As future work it would also be interesting to compare the current results with 3D simulations, in particular regarding the diffusion properties which can be different in 3D \\citep{Jokipii:1993p3406}, and the characteristics of the wave spectra generated for quasi-parallel shocks, both in the low and high Mach number regimes. Finally, a complete statistical study of self-consistently accelerated particles in hybrid and PIC simulations would be important in order to assess the dependence of the diffusion coefficient on shock properties (such as the Mach number and B field inclination angle), and also on the individual particle energies. Such a study could then be used to identify the astrophysical relevant scenarios for the acceleration of cosmic rays."
    },
    "1107/1107.5371_arXiv.txt": {
        "abstract": "We present robust statistical estimates of the accuracy of early-type galaxy stellar masses derived from spectral energy distribution (SED) fitting as functions of various empirical and theoretical assumptions. Using large samples consisting of 40,000 galaxies from the Sloan Digital Sky Survey, of which 5,000 are also in the UKIRT Infrared Deep Sky Survey, with spectroscopic redshifts in the range 0.05 $\\leq$ z $\\leq$ 0.095, we test the reliability of some commonly used stellar population models and extinction laws for computing stellar masses. Spectroscopic ages (t), metallicities (Z), and extinctions (A) are also computed from fits to SDSS spectra using various population models. These constraints are used in additional tests to estimate the systematic errors in the stellar masses derived from SED fitting, where t, Z, and A are typically left as free parameters. We find reasonable agreement in mass estimates among stellar population models, with variation of the IMF and extinction law yielding systematic biases on the mass of nearly a factor of 2, in agreement with other studies. Removing the near-infrared bands changes the statistical bias in mass by only 0.06 dex, adding uncertainties of 0.1 dex at the 95\\% CL. In contrast, we find that removing an ultraviolet band is more critical, introducing 2? uncertainties of 0.15 dex. Finally, we find that stellar masses are less affected by absence of metallicity and/or dust extinction knowledge. However, there is a definite systematic offset in the mass estimate when the stellar population age is unknown, up to a factor of 2.5 for very old (12 Gyr) stellar populations. We present the stellar masses for our sample, corrected for the measured systematic biases due to photometrically determined ages, finding that age errors produce lower stellar masses by 0.15 dex, with errors of 0.02 dex at the 95\\% CL for the median stellar age subsample. ",
        "introduction": "Our understanding of the formation and evolution of early-type galaxies (ETGs) represents a key ingredient in models of galaxy formation. Here, we consider ETGs to be bulge dominated galaxies with passive spectra in their central regions. Observationally, they are characterized by elliptical isophotes, generally redder colors, and a sharp 4000\\AA~break, corresponding to an accumulation of absorption lines of mainly ionized metals, reflected in their rest frame UV/optical colors. This feature is typical of old stellar populations with little to no ongoing star formation.  At higher redshift, most ETGs can only be studied in integrated light, and interpretation of their photometric and spectroscopic properties requires population synthesis models. These single stellar population (SSP) models usually consist of stars born at the same time with equal initial element compositions, evolved using the isochrone synthesis technique (see \\citealp{CB91} for a review), where stars of different masses follow different evolutionary tracks. SSP models can be combined to produce arbitrary star-formation histories, although most studies restrict themselves to histories with exponentially declining star-formation. Such tau-models are parametrized by the e-folding time of this decline, $\\tau$. The spectral energy distributions (SEDs) from a set of models with various parameters (e.g. initial mass function (IMF), star-formation history, age, metallicity, and extinction) is compared to the photometric or spectroscopic observations to derive a best-fit template. A fundamental parameter derived from such SED fitting is the overall normalization of the model relative to the observations,  which gives the galaxy stellar mass content. Indeed, a measurement of the stellar mass of an ETG is involved in many useful scaling relations, such as the \\emph{size-mass} relation \\citep{size-mass} and \\emph{downsizing} \\citep{downsizing}, where the evolutionary history of ETGs is seen to follow different time scales as a function of their stellar mass content. Stellar mass assembly in galaxies is also used in tests of hierarchical models, such as the evolution of the number density and size of both early and late-type galaxies as a function of redshift \\citep{SDSS, GOODS, COSMOS}. Furthermore, upcoming surveys (e.g. PanSTARRS, LSST, DES) will provide only photometry, so it is crucial to understand how to obtain reliable stellar masses from SED fitting techniques.  The degeneracies among the multiple model parameters which are required to reproduce the observed SEDs of the galaxies,  along with the available photometric bandpasses, determine the uncertainty of the mass estimates. For example, models and observations show that the rest-frame near-infrared (NIR) galaxy flux correlates well with stellar mass \\citep{1998MNRAS.297L..23K, downsizing} due to weak contributions from hot, young stars and dust extinction at these wavelengths. But both different model predictions and/or absence of NIR data in fitting the SED can result in stellar masses which differ by a factor of 2, for both low and high redshift samples \\citep{2006ApJ...652...97V}. Of particular interest in this study, the effects of unknown stellar population age, metallicity, and extinction on the stellar masses derived for elliptical galaxies have not been quantified. It is already well known that the age-metallicity degeneracy \\citep{Worthey} cannot easily be broken with broadband colors alone.  All of these factors either depend on or affect the resulting evolution of the spectral energy distribution, models of which have been generated by, e.g. \\cite{BC03}, \\cite{Maraston05}, and \\cite{PEGASE2}. Many phases in stellar evolution are still not well understood, but one key ingredient in these models, the thermally-pulsating asymptotic giant branch (TP-AGB) phase, is particularly influential for determining the galaxy stellar mass for a range of dominant stellar population ages. Light from these stars largely influences the integrated brightness of the NIR continuum, the effects of which were recently highlighted by \\cite{Maraston06} and \\cite{TP-AGB}, who found systematic differences of factors of $\\sim$2 in mass and age over a large range in redshift. This example has shown us the effects of one sensitive parameter in the SED fitting process and demonstrates the need to further test the dependence of the stellar mass estimate on model parameters and the use of common ground-based observation filters.  Several authors \\citep[e.g.][]{2009ApJ...699..486C, IMF_ETGs, 2011arXiv1103.0259R} have studied some of these effects on the photometrically derived stellar mass, quantifying systematic offsets between masses calculated against different model parameters. However, they all suffer from a lack of spectroscopic data which can provide independent constraints on many of the galaxy properties. \\cite{2009MNRAS.394..774L} used a synthetic catalog of intermediate redshift ($1 \\leq z \\leq 2$) ETGs to test the dependence of stellar population models and assumed model parameters on stellar mass. However, they compare a synthetic catalog to only 125 galaxies observed in the GOODS field, using a different SED fitting code. \\cite{COSMOS} and \\cite{2010MNRAS.404.2087B} test the accuracy of large surveys ($\\sim$10,000 and $\\sim$200,000 galaxies, respectively) at estimating photometric stellar masses compared to those from available spectroscopy. However, these projects select mostly galaxies at high redshift and neglect the possible combined effects of assuming various free parameters.  This paper is part of a series examining the global and internal properties of ETGs in the nearby Universe, combining optical and NIR photometry with spectroscopic data. The Spheroids Panchromatic Investigation in Different Environmental Regions (SPIDER) project is described in \\citet[][hereafter Paper I]{PaperI}. The second and third papers of the series present a thorough analysis of the optical+NIR scaling relations of ETGs \\citep{PaperII, PaperIII}, while in Paper IV \\citep{PaperIV} we have analyzed the optical+NIR internal color gradients of ETGs. Here we present an extensive comparison of stellar mass estimates obtained for low redshift ETGs by varying several observational and model parameters (presented individually in Section \\ref{sec_results}). The paper is organized as follows. Section \\ref{sec_sample} describes the galaxies, comprised of a carefully selected sample of high S/N, low redshift ETGs. In Section~\\ref{sec_models} we describe the different stellar population models used to compute model galaxy SEDs. In Section~\\ref{sec_technique} we provide an overview of the general technique applied to fit the theoretical SEDs to the observed colors and the SDSS galaxy spectra. The reliability of these fits is assessed in Section~\\ref{sec_reliability}, by comparing photometric and spectroscopic fitting results. Our method for determining stellar masses is described in Section~\\ref{sec_method}, while the results of the various model comparisons are discussed in Section \\ref{sec_results}. Throughout the paper, we adopt a cosmology with H$_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We measured systematic errors in galaxy stellar masses due to different ingredients in widely used models from stellar population synthesis techniques. Using a large sample of high S/N, low redshift ETGs, we want to accomplish two primary goals  -- Determine at a high-level of significance, the statistical biases and uncertainties inherent in a stellar mass estimate;  -- Provide robust corrections to a sample of photometrically determined galaxy stellar masses, when spectroscopy is unavailable.  We first presented a comparison of BC03, CB10, and PEGASE.2 stellar mass estimates, finding that masses have dispersions of $\\sim$$0.10-0.15$ dex at the 95\\% CL, with this uncertainty arguably being due to the different treatment of the TP-AGB stellar evolutionary phase. In general, our tests find no statistically significant systematic mass bias at the $2\\sigma$ level due to variations of the stellar population models, extinction laws, and photometric bandpasses. This includes varying scaling bandpasses in the fitting technique as well as the magnitudes used in generating the photometric SED. However, a systematic offset of about $-0.23$ dex (inconsistent with zero at the $2\\sigma$ level) is observed between stellar masses predicted from Chabrier and Salpeter IMFs. The newer TP-AGB calculations by \\cite{2007A&A...469..239M} affecting $K$-band magnitudes also improve the quality of $u$-band fits, where the dispersion in galaxy stellar mass measured with and without $u$-band data is improved by nearly 20\\% from BC03 to CB10 models. It is interesting to note that stellar mass estimates are more consistent with no near-infrared photometry than they are without $u$-band data, at least using BC03 models with this sample of ETGs.   The observed systematic stellar mass biases and dispersions for our optical+NIR and complete sample of ETGs are summarized in Table \\ref{tab_summary}. As expected for this redshift range, using the photo-z produces the largest systematic uncertainties in the stellar mass estimate. We find that a \\cite{Chabrier} IMF produces lower stellar masses than a \\cite{Salpeter} IMF by about 0.227 (0.225) dex for BC03 (CB10) models, in agreement with other studies, but lower by about 0.1 dex than the difference predicted by direct integration of the formulae over the mass range 0.1$-$100~$\\mathcal{M}_{\\sun}$ (Appendix~\\ref{app:imf}). For the optical+NIR sample, this difference falls within the 3$\\sigma$ confidence limits of the observational estimate (i.e. 0.23 +0.15/-0.10 for BC03), while for the complete sample, it is within the 4$\\sigma$ confidence limits for BC03 and highly inconsistent with zero ($>4\\sigma$ deviation) for CB10.  An important part of this work is the investigation of systematic effects on the photometrically determined stellar mass when spectroscopic constraints on age, metallicity, and extinction are unavailable, and to provide statistically robust stellar mass corrections as a function of these photometric parameters for low redshift early-type galaxies. We used 5,068 objects with $ugrizYJHK$ photometry and spectroscopy to achieve this goal. Offsets from extinction constraints are the smallest of these systematic effects; this implies that the SED fitting is less sensitive to the coarse extinction grid employed by stellar population models when estimating the stellar mass. We find a negligible trend with $A_{V}$ in stellar mass difference for the optical+NIR sample. However, there is about a 0.15 dex bias towards higher masses for all extinctions when $A_{V}$ is unknown, using M09/CB10 models. The dispersion in these mass differences is about $\\pm0.5$ dex (consistent with zero), slightly lower for BC03 models, yet larger than uncertainties introduced by unknown age/metallicity, mainly because of the smaller number of galaxies in a given extinction bin. Fitting broadband photometry seems to produce lower interstellar extinction than spectral fitting for our sample of ETGs, as shown in the {\\it right} panels in Figure \\ref{fig_starlight_v_lephare_all_delta}. This would lower the $0.4 \\times m_{\\lambda^{\\prime}}$ term in Equation \\ref{eqn_m} and produce a higher stellar mass estimate. We comment that for general surveys, it might be more appropriate to fit an SED with one extinction law \\citep[e.g.][]{Calzetti}, and then refit with a more appropriate law once an approximate age is established.  The positive trend in $\\Delta ( \\log \\mathcal{M} )$ with increasing metallicity is simply a result of the inverse trend shown in the {\\it middle} panels of Figure \\ref{fig_starlight_v_lephare_all_delta}. That is, spectral fits on average produce lower metallicities than broadband photometric SED fitting for these objects; and for any given age, a lower metallicity SSP has a lower mass. This is merely a  reflection of the failure of such a coarse metallicity grid to constrain the model in the SED fitting, as expected. \\cite{1987AJ.....94..899D} show that the size of the 4000\\AA~break in spheroidal galaxies is quite insensitive to metallicity, so even provided a wider range of stellar population metallicities, our result should remain unchanged in this regard for (old) ETGs. We chose a larger grid of stellar ages for M09/BC03/CB10 models, ranging from $0.5-12.5$ Gyr to constrain age in the $\\chi^{2}$-minimization. We measure remarkably lower stellar masses when age is unknown for all young/old populations (except $\\lesssim4$ Gyr for CB10 models), with uncertainties of around $0.2-0.3$ dex. Indeed, for our low redshift ETGs with older stellar populations, we would be underestimating the stellar mass content by a factor of 2 (for galaxies with $t\\sim10$ Gyr from BC03/CB10 models). We can argue that the increasing trend with stellar age in Figure \\ref{fixed_histograms_opticalNIR} is directly due to the overall underestimate in photometrically determined age for all models (see {\\it left} panels of Figure \\ref{fig_starlight_v_lephare_all_delta}), whereas we argue in Section \\ref{sec_age}, a difference of $\\sim$0.7 dex in the best fit age can yield a stellar mass difference of $\\sim$0.4 dex, explained by varying $\\mathcal{M}_{*}/\\mathcal{L}$ ratios between different stellar age SEDs. Notice that unknown ages, resulting from a mismatch between the true star formation history of a galaxy and the simplified form used in SED fits (e.g. an exponentially declining SFR), can also affect significantly the stellar mass determination of high-redshift  galaxies, as discussed, e.g., by \\cite{2007A&A...474..443P} and \\cite{2010ApJ...725.1644L}.  In conclusion, we find uncertainties in the galaxy stellar mass due to different stellar population models, IMFs, and extinction laws that are much less than uncertainties introduced by using the photo-z over this small redshift range. More notable are offsets measured between Chabrier/Salpeter IMFs, Cardelli/CF2000 extinction laws, and changes in TP-AGB evolutionary prescriptions -- between CB10/PEGASE.2 models. The discrepancies in ages, metallicities, and extinctions between spectroscopic and photometric measurement techniques propagate to uncertainties in the stellar mass estimates, yielding notably lower masses when stellar age is unknown. Given the coarse grid of stellar metallicities, the age-metallicity degeneracy is likely to contribute to the large discrepancies in spectroscopic and photometric measurements. Finally, we proposed a \\emph{mass correction} to our sample that incorporates these systematic offsets as a function of their photometrically determined age and metallicity, finding that the sample median mass increases by a factor of roughly 1.3. We would like to comment that these results apply to optical and near-infrared broadband photometry of low redshift ETGs. Detailed quantification of the systematic errors involved in SED fitting have produced results consistent with what other studies have found and have allowed us to obtain consistent stellar mass estimates for this sample.  We have used data from the 4th data release of the UKIDSS survey, which is described in detail in \\citet{War07}. The UKIDSS project is defined in~\\citet{Law07}. UKIDSS uses the UKIRT Wide Field Camera (WFCAM; \\cite{Casali07}). The photometric system is described in \\cite{Hewett}, and the calibration is described in \\cite{Hodgkin09}. The pipeline processing and science archive are described in Irwin et al. (2011, in prep) and \\cite{Hambly08}. 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 \\url{http://www.sdss.org}. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions  are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. The authors would also like to thank Olivier Ilbert and Stephane Arnouts for help with using $LePhare$.  \\begin{sidewaystable} \\centering \\caption{Summary of results: \\emph{optical+NIR}. \\emph{Offset} is the median of the $\\Delta ( \\log \\mathcal{M} )$ histograms and $2\\sigma$ is the $\\pm$2-sigma uncertainty on the median. \\emph{Offset} is provided for the fixed$-$free parameter tests (Figure \\ref{fixed_histograms_opticalNIR}) at the median spectroscopic values of $t$, Z, and $A_{V}$ for BC03 and CB10 models. Unless otherwise noted, these tests assume a spectroscopic redshift, Chabrier IMF, and Cardelli extinction law. Note that galaxies fit with the median spectroscopic age of 4.0 Gyr (CB10) suffered catastrophic fitting failures, and so the offset and uncertainties are reported here for an age of 4.5 Gyr.} \\begin{center} \\begin{tabular}{cccccccc}  \\hline \\hline & \\multicolumn{2}{c}{BC03} & \\multicolumn{2}{c}{CB10} \\\\ & Offset & $2\\sigma$ & Offset & $2\\sigma$ & Note \\\\ \\hline spectro-z $-$ photo-z & 0.0128 & +0.804/-0.500 & 0.0555 & +0.771/-0.524 & $0.000 \\leq z_{p} \\leq 0.300$, $9 < \\log \\mathcal{M}/\\mathcal{M}_{\\sun} < 12$ \\\\ BC03 $-$ CB10 & -4.64E-3 & +0.118/-0.163 & -4.64E-3 & +0.118/-0.163 & STELIB spectral library \\\\ BC03 (CB10) $-$ PEGASE.2 & 0.0962 & +0.0837/-0.112 & 0.0858 & +0.0934/-0.132 & --- \\\\ \\hline Chabrier $-$ Salpeter & -0.227 & +0.0855/-0.0323 & -0.225 & +0.0441/-0.0281 & STELIB spectral library \\\\ \\hline Cardelli $-$ Calzetti & -4.50E-3 & +0.0711/-0.140 & -1.00E-4 & +0.109/-0.0375 & Cardelli/Calzetti from {\\it LePhare} code \\\\ Cardelli $-$ CF2000 & -0.104 & +0.142/-0.170 & -0.0680 & +0.156/-0.131 & $\\tau_{V}=1.0$, $\\mu=0.3$ (in GALAXEV code) \\\\ \\hline $u-$ no-$u$ & -0.0257 & +0.131/-0.151 & -6.30E-3 & +0.115/-0.121 & --- \\\\ $ugrizYJHK-ugriz$ & -0.0566 & +0.0868/-0.125 & -0.0419 & +0.147/-0.108 & --- \\\\ \\hline Fixed$-$Free & & & \\\\ $t$ & 0.0833 & +0.108/-0.144 & 0.0590 & +0.284/-0.201 & $t_{med}=3.25~(4.50)$ Gyr, BC03 (CB10) \\\\ $Z$ & 0.0587 & +0.191/-0.0442 & -1.50E-3 & +0.0273/-0.151 & [Fe/H]$_{med}=0.0932~(0.0932)$, BC03 (CB10) \\\\ $A_{V}$ & 0.0276 & +0.280/-0.147 & -0.164 & +0.554/-0.488 & $A_{V,med}=0.12~(0.10)$, BC03 (CB10) \\\\ \\hline  \\end{tabular} \\end{center} \\label{tab_summary} \\end{sidewaystable}  \\clearpage  \\begin{sidewaystable} \\centering \\caption{Summary of results: \\emph{complete}. \\emph{Offset} is the median of the $\\Delta ( \\log \\mathcal{M} )$ histograms and $2\\sigma$ is the $\\pm$2-sigma uncertainty on the median. \\emph{Offset} is provided for the fixed$-$free parameter tests (Figure \\ref{fixed_histograms_opticalNIR}) at the median spectroscopic values of $t$, Z, and $A_{V}$ for BC03 and CB10 models. Unless otherwise noted, these tests assume a spectroscopic redshift, Chabrier IMF, and Cardelli extinction law. Again, the reader is cautioned on use of the results in this table, per the reasons outlined in Section \\ref{sec_results}.} \\begin{center} \\begin{tabular}{cccccccc}  \\hline \\hline & \\multicolumn{2}{c}{BC03} & \\multicolumn{2}{c}{CB10} \\\\ & Offset & $2\\sigma$ & Offset & $2\\sigma$ & Note \\\\ \\hline spectro-z $-$ photo-z & 0.0281 & +0.780/-0.330 & 0.0185 & +0.772/-0.316 & $0.000 \\leq z_{p} \\leq 0.300$, $9 < \\log \\mathcal{M}/\\mathcal{M}_{\\sun} < 12$ \\\\ BC03 $-$ CB10 & 0.0131 & +0.0424/-0.0326 & 0.0131 & +0.0424/-0.0326 & STELIB spectral library \\\\ BC03 (CB10) $-$ PEGASE.2 & 0.0871 & +0.0516/-0.0908 & 0.0802 & +0.0550/-0.151 & --- \\\\ \\hline Chabrier $-$ Salpeter & -0.233 & +0.0144/-0.0140 & -0.231 & +0.0173/-0.0151 & STELIB spectral library \\\\ \\hline Cardelli $-$ Calzetti & -0.0193 & +0.0777/-0.0386 & 5.10E-3 & +0.0578/-0.0718 & Cardelli/Calzetti from {\\it LePhare} code \\\\ Cardelli $-$ CF2000 & -0.0693 & +0.0918/-0.0501 & -0.0225 & +0.0941/-0.0433 & $\\tau_{V}=1.0$, $\\mu=0.3$ (in GALAXEV code) \\\\ \\hline $u-$ no-$u$ & -0.0128 & +0.0410/-0.102 & -0.0166 & +0.0571/-0.187 & --- \\\\ \\hline  \\end{tabular} \\end{center} \\label{tab_summary_complete} \\end{sidewaystable}  \\clearpage  \\appendix"
    },
    "1107/1107.1552_arXiv.txt": {
        "abstract": "Active Galactic Nuclei (AGN), such as Seyfert galaxies, quasars, etc., show light variations in all wavelength bands, with various amplitude and in many time scales. The variations usually look erratic, not periodic nor purely random. Many of these objects also show lognormal flux distribution and RMS - flux relation and power law frequency distribution. So far, the lognormal flux distribution of black hole objects is only observational facts without  satisfactory explanation about the physical mechanism producing such distribution in the accretion disk. One of the most promising models based on cellular automaton mechanism has been successful in reproducing PSD (Power Spectral Density) of the observed objects but could not reproduce lognormal flux distribution. Such distribution requires the existence of underlying multiplicative process while the existing SOC models are based on additive processes. A modified SOC model based on  cellular automaton mechanism for producing lognormal flux distribution is presented in this paper. The idea is that the energy released in the avalanche and diffusion in the accretion disk is not entirely emitted instantaneously as in the original cellular automaton model. Some part of the energy is kept in the disk and thus increase its energy content so that the next avalanche will be in higher energy condition and will release more energy. The later an avalanche occurs, the more amount of energy is emitted to the observers. This can provide multiplicative effects to the flux and produces lognormal flux distribution. ",
        "introduction": "Black hole objects, in a wide range of mass from  Active Galactic Nuclei (AGNs) to black hole X-ray binaries (BHBs), generally show light fluctuation which is not periodic, nor purely random. This kind of variability is found in all wavelengths and in a wide range of time scales. It is one of the keys for understanding the nature of these objects.  Long term monitoring of AGN light curves still fail to define clearly the upper limit of variability time scale, constraining our ability to extract information about disk structure from variability studies \\citep{hawkins06}. For instance, twenty eight years of monitoring program of about 1500 quasars by \\cite{hawkins96} still did not reach the common upper limit of quasars variability time scale. However, analysis of quasars long term light curve in quasars rest frame using Fourier power spectrum yield a break frequency. \\cite{hawkins07} interpreted it as the 11 year characteristic time scale of quasars.  Several attempts have been done to reproduce such variability, for example, \\cite{Terrell72} introduced a shot noise model, which was developed further by several authors \\citep[e.g.][]{{Lehto89, Lochner91}}. This model could reproduce the observed light curve and power spectral density (PSD) of the AGNs by superposition of the shots that occur randomly. In addition, the shots have identical time profiles with exponential decays. However, there is no clear physical explanation for the creation of the shots \\citep{Takeuchi97}. Other models of variability that have been proposed so far, include starburst, disk instability and microlensing models \\citep[][and references therein]{hawkins04}.  \\cite{mineshige94a, mineshige94b} have tried to simulate the fluctuation in the light curve of black hole objects under the notion of self organized criticality (SOC) by using a cellular automaton rule. SOC is a continuous critical state reached spontaneously by a system receiving energy and/or matter from outside, independent of its initial condition.  The concept of SOC was first introduced by Bak et al. (1998), which used sand pile system as an illustration. When grains of sand are dropped continuously in a surface, it will accumulate and create a peak which is higher and higher. When the peak is high enough and the slope of some part of the pile reach a critical value, then an additional falling grain of sand will trigger another grain to fall one after the other and creating an avalanche. After an avalanche the slope of the pile will be shallower. Additional sand drop will make the surface of the pile steeper and steeper, finally reach a critical value and trigger another avalanche, and so on. The sand pile will always be close to the critical state which may trigger avalanches. If the energy release of this sand pile system is recorded in a time series, it will produce $1/f$-like variability in its PSD.  In their model, \\cite{mineshige94a} and Mineshige et al. (1994b) divided the two-dimensional disk in radial and azimuthal directions, turned into cells. The disk receives mass input from a randomly chosen outermost cell and finally loses mass through the innermost ring due to accretion onto its central black hole. Similar to the sand pile model, in the accretion disk the mass is transfered inward through avalanche and diffusion processes that may happen in any cell under some special conditions. These processes convert gravitational potential energy into heat, eventually radiated from the disk. By using this model, they could successfully simulate the light fluctuation and its PSD. The PSD showed a $1/f$ tendency with power law index similar to those observed from some black hole objects.  However, later observations of some AGNs revealed that some of them show lognormal X-ray flux distribution \\citep[e.g.][]{Gaskell04} and follow linear RMS-flux relation \\citep[e.g.][]{uttley05}. In this case RMS is the root mean square of the flux and the flux means the average flux for the chosen bin. These properties are also found in some galactic accreting source, such as Cyg X-1 and SAX J1808.4-3658 \\citep{uttley01, Gleis04}. The SOC model introduced by \\cite{mineshige94a, mineshige94b} could not explain such phenomena \\citep[see e.g.][]{Gaskell04}, because, multiplicative processes are necessary to yield lognormal flux distribution, whereas the underlying process in the original SOC model is additive. What kind of physical processes in the accretion disk can yield lognormal flux distribution? Some mathematical models have been introduced \\citep[see e.g.][]{Gaskell04, uttley05}, but without a clear explanation on the physical mechanisms.  In this work, we introduce an idea with reasonable physical mechanism which has multiplicative effect, in an attempt to obtain a lognormal flux distribution. We assume that the accretion flow is dominated by advective cooling instead of radiative cooling. This mechanism was called advection-dominated accretion flow (ADAF) \\citep{Ichimaru77, Abra88, Kato08}.  In the present work, we adopt the idea that the energy released by the viscous process is not efficiently emitted locally which could act as a multiplicative effect that yield lognormal flux distribution. It is important to note that \\cite{Takeuchi97} created a numerical simulation of hydrodynamical model of advection dominated accretion disk with critical behavior. They aimed to reproduce $1/f$-like X-ray fluctuation with a physical mechanism focusing on the Cyg X-1 type fluctuation. We, on the other hand, aimed to address the lognormal flux distribution problem that was claimed to be unable to be reproduced by a cellular automaton model \\citep[see e.g.][]{Gaskell04}.  The plan of this paper is as follows: We will describe the modified cellular automaton model in the second section in which we will also briefly review the lognormal distribution and the original cellular automaton model. In section three, we present the simulation details. The results and analysis will be presented in section 4. Section five is devoted for discussion and conclusion. ",
        "conclusions": "The characteristics of simulated light curve which appear in PSD, structure function, flux distribution, and RMS - flux relation are similar with those derived from observation. The slope of the PSD from our simulation is about 1.9, this is in the range of the PSD slope of AGNs \\citep[see for example][]{klis95, starling04, uttley02}.  The structure function can provide some sort of information about the variability time scale. The structure function yielded goes along with the observation \\citep[see for example][]{vries03} except the saturation part which provided the hint on longest variability time scale. We can see saturation part in our simulation, but not in the data of 30 years monitoring data of quasars reported by \\cite{vries03}. This indicates that the longest variability time scale of quasars is more than 30 years.  The flux distribution of simulated light curve is lognormal which is agree with observation. \\cite{Gaskell04} found a lognormal flux distribution in an extreme narrow-line Seyfert 1 Galaxy, and argued that the lognormal intensity distribution of AGN is against the argument that the AGN variations being driven by self organized criticality. That is correct for the old version of SOC model. We, however successfully made some modification of the energy emission mechanism by combining ADAF and cellular automaton mechanism which can yield lognormal distribution, while keeping the other characteristics goes along with the observation.  This result can revive the SOC model for accretion disk which begin to fade away and more elaborate work based on SOC model is necessary to develop a better model with a hope to finally simulate the real situation in accretion disks. The model developed here is still rough and applied to the accretion disk in a not very detail manner, it's just to show that SOC model can still yield lognormal distribution and RMS-flux relation."
    },
    "1107/1107.4007_arXiv.txt": {
        "abstract": "We report the detection, with the CANGAROO-III imaging atmospheric Cherenkov telescope array, of a very high energy gamma-ray signal from the unidentified gamma-ray source HESS~J1614$-$518, which was discovered in the H.E.S.S.\\ Galactic plane survey.  Diffuse gamma-ray emission was detected above 760\\,GeV at the 8.9~$\\sigma$ level during an effective exposure of 54\\,hr from 2008 May to August. The spectrum can be represented by a power-law: $(8.2\\pm2.2_{stat}\\pm2.5_{sys})\\times10^{-12}\\times(E/1~$TeV$)^{-\\gamma}$ cm$^{-2}$~s$^{-1}$~TeV$^{-1}$ with a photon index $\\gamma$ of $2.4\\pm0.3_{stat}\\pm0.2_{sys}$, which is compatible with that of the H.E.S.S.\\ observations.  By combining our result with multi-wavelength data, we discuss the possible counterparts for HESS~J1614$-$518 and consider radiation mechanisms based on hadronic and leptonic processes for a supernova remnant, stellar winds from massive stars, and a pulsar wind nebula. Although a leptonic origin from a pulsar wind nebula driven by an unknown pulsar remains possible, hadronic-origin emission from an unknown supernova remnant is preferred. ",
        "introduction": "Recent progress with Imaging Air Cherenkov Telescopes (IACTs) is enabling the exploration of sites of cosmic-ray acceleration in our galaxy. Very high energy (VHE) gamma rays are produced by the decay of neutral pions which arise from interactions between the accelerated protons and interstellar matter, or by Inverse Compton (IC) scattering and Bremsstrahlung of high-energy electrons.  For example, VHE gamma rays have been detected from young supernova remnants (SNRs) such as RX J1713.7$-$3946 \\citep{Muraishi_2000,can_1713,hess_1713}, RX J0852.0$-$4622 \\citep{can_0852,hess_0852,enomoto2}, and RCW86 \\citep{Watanabe_2003,hess_rcw86}, which show possible evidence of cosmic-ray acceleration \\citep[e.g.,][]{can_1713,2005ApJ...624L..37M,Uchiyama_1713,tanaka_1713,yamazaki_2009}. In addition, detections of VHE gamma rays from pulsar wind nebulae (PWNe) such as the Crab nebula \\citep{Weekes} and Vela X nebula \\citep{velax,enomoto} have shown that PWNe also play an important role in particle acceleration in the Galaxy.  Recently, VHE gamma-ray emission related to massive stars such as the X-ray binary systems PSR~B1259$-$63 \\citep{1259} and LS~I~+61 303 \\citep{LSI}, the young open stellar clusters Cyg~OB2 \\citep{2032}, Westerlund~2 \\citep{wes2}, and Westerlund~1 \\citep{2010ASPC..422..265O} have also been reported. Moreover, the Galactic plane survey performed by the H.E.S.S.\\ observatory \\citep{hess,hess2} discovered seventeen unidentified VHE gamma-ray sources, including HESS~J1614$-$518. Today, unidentified sources are the largest class of the 123 discovered VHE gamma-ray sources, most of which are located in the Galactic plane (e.g., \\citet{2005A&A...431..197A,milagro,hess_unid}; the TeVCAT catalog, \\citet{TeV_catalog}, is a useful up-to-date, on-line resource).  In general, the lack of non-thermal electromagnetic radiation from the radio to the X-ray bands may be evidence of hadron acceleration because the IC scenario requires a lower magnetic field than the typical interstellar magnetic field intensity of a few $\\mu$G.  Revealing the possible radiation mechanism(s) of each unidentified source is therefore important for identifying the origin(s) of cosmic rays.  H.E.S.S.\\ reported that HESS~J1614$-$518 had a high flux level, 25\\% of the Crab nebula, above 200\\,GeV with a photon index of 2.4 and an elliptical morphology with a semi-major axis of $14\\pm1$~arcmin and a semi-minor axis of $9\\pm1$~arcmin \\citep{hess2}. The peak position has an offset of 8.7~arcmin to the north-east from the central position. \\citet{landi2} and \\citet{rowell} pointed out that HESS J1614$-$518 may be associated with the 40~Myr-old young open star cluster Pismis~22 \\citep{piatti_pismis} which is located within the VHE gamma-ray emission region at a distance of 1.0$\\pm$0.4\\,kpc and has sufficient luminosity to produce the observed gamma-ray luminosity, assuming 20\\% energy conversion from the stellar winds of ten B-type stars.  However, there are several issues in identifying HESS~J1614$-$518 with Pismis~22 since the size of Pismis~22, 2.0\\,arcmin in diameter, is one order of magnitude smaller than the VHE gamma-ray emission size and the location has a 12\\,arcmin offset from the VHE gamma-ray emission peak.  In addition, there has been no detailed discussion of the radiation mechanism.  The X-ray satellite {\\it Suzaku} observed this region with the X-ray imaging spectrometer (XIS) and found three X-ray sources in their follow-up observation in 2006 \\citep{matsumoto}.  One of these, called Suzaku source~A, is located very close to the VHE gamma-ray peak position with an offset of 0.8\\,arcmin. The spectrum is well fitted by a single power-law model with a photon index of $1.73^{+0.33}_{-0.30}$ and a hydrogen equivalent column density of $1.21^{+0.50}_{-0.41}\\times10^{22}$ cm$^{-2}$.  The distance to Suzaku source~A is approximately 10~kpc, which was derived from the hydrogen equivalent column density using the total Galactic H\\,{\\sc i} column density toward HESS~J1614$-$518 of $\\sim2.2\\times10^{22}$~cm$^{-2}$ \\citep{Dickey}.  The size of the X-ray emission region is slightly larger than the {\\it Suzaku} Point Spread Function (PSF) of 1.8\\,arcmin and smaller than the size of the VHE gamma-ray region.  This difference is seen in PWNe such as HESS~J1825$-$137 \\citep{hess3,uchiyama_1825} and Vela X \\citep{velax_X,velax} and could be explained by the difference between the synchrotron cooling time of the electrons that radiate X-rays and those that produce TeV gamma rays. Electrons with an energy of 100\\,TeV radiate X-rays and immediately lose their energy by synchrotron cooling (e.g., the energy-loss timescale is $\\sim10^{2}$~yr assuming a magnetic field of 20~$\\mu$G \\citep{Sturner_1997}), while electrons with an energy of 1\\,TeV, which are responsible for the VHE gamma ray emission through IC scattering, are more slowly cooled by synchrotron radiation (e.g., the energy-loss timescale is $\\sim10^{4}$~yr assuming the same parameters as above) and can travel further from their source.  However, the ratio between the observed VHE gamma-ray and X-ray fluxes, F(1--10 TeV) / F(2--10 keV) of $\\sim$34, is much larger than those of known PWNe --- $2.6\\times10^{-3}$, 0.7, and 1.5 for the Crab, MSH~15-52, and Vela X, respectively \\citep{Gaensler_1999, x_crab,Gaensler_2002, Dodson_2003, hess_crab,Aharonian_2005, velax, Manzali_2007, Nakamori_2008}. Nevertheless, recent studies of HESS~J1640$-$465 \\citep{funk} and HESS~J1804$-$216 \\citep{Higashi_2008} claim that this large ratio can be explained by a time-evolving electron injection model, in which the number of electrons injected into space by the pulsar decreases proportionally to the spin-down of the pulsar.  On the other hand, this large ratio is also expected in an old SNR with an age of $\\sim10^{5}$~yr, because of the difference between the cooling times of electrons and protons \\citep{yamazaki_old}.  We therefore discuss both a PWN scenario and an SNR scenario in this paper.  Suzaku source~B is positioned towards the center of HESS~J1614$-$518 and is coincident with the position of Pismis~22. Since the hydrogen equivalent column density derived from the Suzaku spectrum is $(1.1\\pm0.21)\\times10^{22}$ cm$^{-2}$, which is comparable with that of Suzaku source~A, Suzaku source~B may lie at a similar distance as for Suzaku source A. This source has a non-thermal X-ray emission with a photon index of $3.19\\pm0.32$. This soft index and X-ray luminosities of $7.7\\times10^{34}$~ergs~s$^{-1}$ and $4.5\\times10^{35}$~ergs~s$^{-1}$ in the 2$-$10~keV and 0.5$-$10~keV ranges, respectively, assuming a distance of 10~kpc, are typical values for an anomalous X-ray pulsar (AXP) \\citep{AXP_01,AXP_02}. The possible existence of the AXP also suggests this source as an SNR, since AXPs are usually associated with SNRs, e.g., 1E~2259+586 with CTB~109 \\citep{AXP_01} and 1E~1841$-$045 with Kes~73 \\citep{AXP_1841}. Since the position of this source is coincident with Pismis~22, there is a possibility that this emission originates from the stellar winds from the stellar cluster. Non-thermal X-ray emission from a stellar cluster was reported from Westerlund~1 \\citep{Muno_2006}, and TeV gamma rays were recently detected from this object \\citep{Ohm_2009}. However, this positional correlation may be only a chance coincidence since the estimated distances to Suzaku source~B and Pismis~22 are different by an order of magnitude. Although Suzaku source~B might be marginally extended, it is difficult to quantitatively estimate the spatial extension with the Suzaku PSF of 1.8~arcmin. If Suzaku source B is actually extended, additional scenarios besides an SNR could be considered, e.g., a PWN from a pulsar/AXP as discussed in \\citet{matsumoto}, or emission from the unresolved hot stars in Pismis 22.  The other source, Suzaku source~C, is a late B-type star as described in \\citet{matsumoto}, and thus is not a possible counterpart of HESS J1614$-$518.  {\\it Swift} observed this region with the X-ray telescope (XRT) and found six X-ray sources (hereafter Swift sources~1 to 6) \\citep{landi1,landi2}.  All these sources were point-like and no diffuse emission was found.  Two sources, Swift sources~1 and 4, are located within the field of view (FOV) of the {\\it Suzaku} observation.  Swift source~1 is located close to Pismis~22 with an offset of 42~arcsec.  This source is also coincident with Suzaku source~B. Swift source~4 is coincident with Suzaku source~C.  Swift sources~1, 2, 3, 5 are probably stars, while the nature of Swift sources~4 and 6 were not identified, probably due to the poor statistics.  Although Suzaku source~A was located in the FOV of the {\\it Swift} XRT, it was not detected with {\\it Swift} probably due to the limited exposure time ($\\sim$1700~s) and/or the small effective area.  The Fermi-LAT collaboration \\citep{fermi_cat} reported the detection of gamma rays in the 100\\,MeV to 100\\,GeV band from 1FGL~J1614.7-5138c positioned 2.7\\,arcmin away from the peak position of the VHE gamma-ray emission.  In the radio band, no counterpart has been found in the HESS~J1614$-$518 region; there is no enhancement in the 843~MHz band, where the rms noise level is $\\sim$2\\,mJy\\,arcmin$^{-2}$ \\citep{radio,Murphy_2007}.  In this paper, we present TeV gamma-ray observations of HESS~J1614$-$518 with the CANGAROO-III telescopes and discuss the possible counterpart and the radiation mechanism by considering multi-wavelength observations. ",
        "conclusions": "The observation of HESS J1614$-$518 with the CANGAROO-III telescopes confirms the VHE gamma-ray emission reported by H.E.S.S. The differential energy spectrum can be fitted with a single power law:$(8.2\\pm2.2_{stat}\\pm2.5_{sys})\\times10^{-12}\\times(E/1 $TeV$)^{-\\gamma}$ cm$^{-2}$s$^{-1}$TeV$^{-1}$ with a photon index $\\gamma$ of $2.4\\pm0.3_{stat}\\pm0.2_{sys}$. We discuss the possible counterparts for this object using the results of observations with {\\it Suzaku} and {\\it Fermi}. For the SNR scenario, a one-zone leptonic model was not able to account for the observed SED. Hadronic models gave a good reproduction of the SED and the typical SNR explosion energy of $\\sim10^{51}$\\,ergs is able to supply the total energy of protons. Since the required number densities of the ambient matter were $n_{p}=1~p$~cm$^{-3}$ and $n_{p}=100~p$~cm$^{-3}$ for a distance to the SNR of 1~kpc and 10~kpc, respectively, detailed molecular observations could determine whether the SNR originated from Pismis~22 (d$\\sim$1~kpc) or a farther distance. As there were also differences in the spectrum of the emission from the secondary electrons, a determination of the spectrum below the infrared band would help determine the likelihood of an SNR origin. For the PWN scenario, the nearby known pulsars are not responsible since the spin-down powers are insufficient to produce the observed TeV gamma-ray luminosity. Further observations to search a pulsar are necessary to investigate the PWN scenario. For the stellar wind scenario, Pismis~22 was required to contain two O-type stars through its entire age from energetics considerations. However, the required number density of the ambient matter of $n_{p}=100~p$~cm$^{-3}$ may not be consistent with the results of the NANTEN observations.  To identify HESS~J1614$-$518, more detailed multiwavelength observations are required. To discuss the stellar wind origin in more detail, a determination of the number of OB stars is necessary. For the SNR scenario, the ultra-high energy resolution of the SXS onboard {\\it Astro-H} could detect line emissions with a high abundance of $\\alpha$-elements compared to that of iron, which would indicate that HESS~J1614$-$518 is an SNR. To show that the Suzaku source~B is an AXP, supporting the SNR or PWN scenario, a high time-resolution X-ray observation is needed to detect a pulsed signal from the source. More detailed gamma-ray spectroscopy with {\\it Fermi}, and the Cherenkov Telescope Array \\citep{CTA} could determine the origin of the accelerated particles."
    },
    "1107/1107.3405_arXiv.txt": {
        "abstract": "{ Cosmic rays are hampered by the Moon and a deficit in its direction is expected (the so-called \\emph{Moon shadow}). The Moon shadow is an important tool to determine the performance of an air shower array. In fact, the displacement of the shadow center, due to the bending effect of the Geomagnetic field on the propagation of cosmic rays, allows to set the energy scale of the primary particles inducing the showers observed by the detector. The shape of the shadow permits to determine the detector point spread function. The position of the deficit at high energy allows evaluating its pointing accuracy. Here we present the observation of the cosmic ray Moon shadowing effect carried out by the ARGO-YBJ experiment (Yangbajing Cosmic Ray Laboratory, Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$ ) in the multi-TeV energy region with high statistical significance (70 standard deviations). By means of an accurate Monte Carlo simulation of the cosmic rays propagation in the Earth-Moon system we have studied the role of the Geomagnetic field and of the detector point spread function on the observed shadow.} ",
        "introduction": "The analysis of the Moon shadow observed by an air shower array may provide unique information on its performance. At high energies, the Moon shadow would be observed by an ideal detector as a 0.52$^{\\circ}$ wide circular deficit of events, centered on the Moon position. The actual shape of the deficit as reconstructed by the detector allows the determination of the angular resolution while the position of the deficit allows the evaluation of the absolute pointing accuracy. In addition, charged particles are deflected by the Geomagnetic field (GMF) by an angle depending on the energy. As a consequence, the observation of the displacement of the Moon shadow at low rigidities can be used to determine the relation between the shower size and the primary energy. The same shadowing effect can be observed in the direction of the Sun but the interpretation of this phenomenology is less straightforward \\cite{argo-solar}. Finally, the Moon shadow can be exploited to measure the antiproton content in the primary cosmic rays (CR). In fact, acting the Earth-Moon system as a magnetic spectrometer, paths of primary antiprotons are deflected in the opposite sense with respect to those of the protons in their way to the Earth. This effect has been used to set limits on the $\\bar{p}$ flux at TeV energies not yet accessible to balloon/satellite experiments \\cite{antip}.  In this paper we present the observation of the cosmic ray Moon shadowing effect carried out by the ARGO-YBJ experiment during the period from July 2006 to November 2010. We report on the angular resolution, the pointing accuracy and the energy scale calibration of the detector in the multi-TeV energy region. The results are compared with the predictions of a detailed simulation of cosmic ray propagation in the Earth-Moon system. ",
        "conclusions": "The Moon shadowing effect on cosmic rays has been observed by the ARGO-YBJ experiment in the multi-TeV energy region with a statistical significance greater than 70 standard deviations.  We have estimated the primary energy of the detected showers by measuring the westward displacement as a function of the multiplicity, thus calibrating the relation between shower size and CR energy. The systematic uncertainty in the absolute energy scale is evaluated to be less than 13\\% in the range from 1 to 30 (TeV/Z).  The position of the Moon shadow measured with the ARGO-YBJ experiment turned out to be stable at a level of 0.1$^{\\circ}$ and the angular resolution stable at a level of 10\\%, on a monthly basis. These results make us confident about the detector stability in the long-term observation of gamma-ray sources."
    },
    "1107/1107.4013_arXiv.txt": {
        "abstract": "{This paper presents U\\,Sco nebular spectra collected in the period March-May 2010 after the binary outburst on Jan 28, 2010. The spectra display strong [Ne{\\sc v}] and [Ne{\\sc iii}] lines that can be used to compute the relative abundance of [Ne/O]. The value obtained ([Ne/O]=1.69) is higher than the typical [Ne/O] abundance found in classical novae from CO progenitors and suggests that U\\,Sco has a ONeMg white-dwarf progenitor. It follows that U\\,Sco will not explode as a SN\\,Ia but rather collapse to become a neutron star or a millisecond pulsar.} ",
        "introduction": "The outburst of U\\,Scorpii (U\\,Sco) on Jan 28, 2010 at $V_{\\sc max}=8.05$\\,mag (Munari et al. 2010) is the tenth outburst recorded for this nova and, thanks to the Schaefer monitoring and alert (see e.g. Schaefer et al.\\ 2010), the best-studied of this kind together with the outburst of 1999. The importance of observing U\\,Sco and recurrent novae (RNe) outbursts in general resides in the opportunity to understand the kinematics and the composition of RNe ejecta in relation to both classical novae (CNe) and supernovae (SNe). While observations (Della Valle and Livio 1996) show that RNe are not the main contributors to the class of SN\\,Ia progenitors, numerous theoretical works consider  RNe within the single degenerate scenario as valid progenitors of type Ia SNe (e.g. Hachisu et al.\\ 2000; Hachisu and Kato 2002; Livio 2000 and reference therein; Nomoto et al.\\ 2000; see also Podsiadlowski 2008). To better understand this problem, I collected X-Shooter multi-wavelength medium resolution spectra of U\\,Sco during its late decline from the plateau phase down to quiescence. The goal was to characterize the physical parameters of the ejecta when they are optically thin and compute  abundances, which are a discriminating factor between eroded white dwarfs (WDs) vs. non-eroded WDs.  In this letter, I present nebular spectra of U\\,Sco taken 45, 73, and 104 days after the outburst and discuss the measured [Ne/O] abundance value in the attempt to identify the underlying WD. A more detailed discussion of the U\\,Sco 2010 outburst spectral evolution will be presented in a separate paper (Mason et al. in preparation), which collects a larger sample of spectra. ",
        "conclusions": "The general consensus about SN\\,Ia progenitors is that they originate from CO WDs accreting matter up to the Chandrasekhar limit either via stellar merging (double degenerate scenario where two CO WDs coalesce) or via mass transfer from a less evolved companion (single degenerate scenario). Within the single degenerate scenario the accreting WD ought to be massive (i.e. close to the Chandrasekhar limit), and accreting at a high rate ( $\\geq$10$^{-8}$-10$^{-7}$\\,M$_\\odot$\\,yr$^{-1}$, Nomoto et al. 2007 and reference therein, see also Yaron et al. 2005). Hence, RNe whose frequent outbursts are explained by massive WDs and high mass-transfer rates (e.g. Truran et al. 1988) remain  viable candidates, contrary to what  is suggested by the population synthesis simulations (Di Stefano 2010 and reference therein) and their census (Della Valle and Livio 1996). U\\,Sco, having a WD $>$1.37M$_\\odot$ (Hachisu et al. 2000; Thoroughgood et al. 2001) and accreting at a rate $>10^{-7}$\\,M$_\\odot$\\,yr$^{-1}$ (Hachisu et al. 2000; Matsumoto et al. 2003), seems very likely to explode as a SN\\,Ia. However, it is also general consensus that an accreting ONeMg WD cannot lead to a SN\\,Ia explosion but, eventually, to a core collapse and the formation of a neutron star or a millisecond pulsar (e.g. Nomoto and Kondo 1991). The high U\\,Sco [Ne/O] abundance (when compared to the solar values of $-$0.76) is intriguing both because it points to dredged-up material and a possibly eroded WD, and because it provides information about the composition/nature of the underlying WD. Dredged-up material is typically associated with an eroded WD because, according to thermo-nuclear (TNR) computations (e.g. Prialnik and Livio 1995), anything that is mixed within the H-layer where the TNR ignites is ejected into the circumbinary space. If at any outburst the mass of the ejecta matches or exceeds the accreted mass, then the WD (independent of its composition) cannot increase in mass and reach the Chandrasekhar limit. On the other hand, if the primary star is an ONeMg WD, it will never explode as a SN\\,Ia even in case of mass gain up to the Chandrasekhar limit of 1.4M$_\\odot$. Hence, it has become critical to establish the composition of the accreting primary for U\\,Sco because despite its WD mass, high accretion rate and small ejecta mass (a few $10^{-7}$\\,M$_\\odot$, Iijima 1999; Anupama and Dewangan 2000, but see also Diaz et al. 2010), it might not explode as a SN\\,Ia but undergo core collapse at the very most.  \\begin{table} \\scriptsize \\centering \\caption{Ne abundances for a sample of classical novae. The [Ne/O] abundances have been computed assuming the solar values in Asplund et al. (2009). } \\begin{tabular}{lccc} \\hline nova   & WD & [Ne/O] & ref. \\\\ \\hline V693 CrA  & ONeMg & +1.21 &  Vanlandingham et al. (1999) \\\\ V1370 Aql  &  ONeMg  &  +1.76 & Livio \\& Truran (1994) \\\\ V1974 Cyg  & ONeMg &  +0.47 & Vanlandingham et al. (2005) \\\\ V382 Vel   & ONeMg &  +0.66 & Shore et al. (2003) \\\\ V4160 Sgr  & ONeMg &  +0.69 & Schwarz et al. (2007a) \\\\ V838 Her  & ONeMg &  +1.61 & Schwarz et al. (2007a) \\\\ QU Vul    & ONeMg &  +0.92 &  Schwarz (2002)\\\\ V1065 Cen   & ONeMg & +0.74 & Helton et al. (2010) \\\\ LMC 1990 N.1 & ONeMg  &  +0.45 & Vanlandingham et al. (1999) \\\\ U Sco  & ONeMg & +1.69 & this paper \\\\ RR Pic    & ? & +0.94  & Livio \\& Truran (1994)\\\\ V977 Sco   & ? & +0.60 & Livio \\& Truran (1994) \\\\ V1186 Sco   & CO & $-$0.42 & Schwarz et al. (2007b) \\\\ LMC 1991 & CO & $-$1.65 & Schwarz et al. (2001) \\\\ QV Vul  &  CO & $-$0.96 & Livio \\& Truran (1994)\\\\ HR Del   & CO & $-$0.53 & Livio \\& Truran (1994)\\\\ V1500 Cyg    & CO & $-$0.069 & Livio \\& Truran (1994) \\\\ V1688 Cyg   &  CO  & $-$0.62 & Livio \\& Truran (1994)\\\\ GQ Mus  & CO & $-$1.09 & Livio \\& Truran (1994) \\\\ PW Vul  & CO &  $-$0.67 & Livio \\& Truran (1994)\\\\ V842 Cen   & CO  & $-$0.86 & Livio \\& Truran (1994) \\\\ V827 Her   &  CO  & $-$0.73 & Livio \\& Truran (1994) \\\\ V2214 Oph   &  CO  & +0.11 & Livio \\& Truran (1994) \\\\ V443 Sct    & CO  & $-$1.04 & Livio \\& Truran (1994) \\\\ \\hline \\end{tabular} \\end{table}  Livio and Truran (1994) cautioned observers about identifying ONeMg WD progenitors in CN ejecta that show Ne emission lines in their optical spectra. They showed that moderate Ne abundances (with respect to solar) can be explained either by abundance uncertainties or dredged-up material from the underlying CO WD, or with the breakout of the CNO cycle under special conditions. The authors identified a group of ``true ONeMg WDs'' in those novae that showed extreme enrichment of Ne and heavier elements. Table~2 lists the novae used by Livio and Truran (1994) for their analysis, as well as a number of CNe whose WD and abundances have been determined via photo-ionization modeling of UV and optical observations, simultaneously, by Schwarz and collaborators. % The table reports the [Ne/O] abundances for each nova as well as the U\\,Sco April\\,11 relative abundance derived in this paper. From Table~2 it is evident that the CNe hosting a CO WD are characterized by [Ne/O] abundances $\\leq$0, while those possessing an ONeMg WD have [Ne/O]$>$0. This is also shown in Fig.3, which plots the distribution of CNe as a function of their [Ne/O] abundance. The shaded histogram of Fig.3 corresponds to the ``fiducial sample'' of ONeMg CN white dwarfs; while the black areas correspond to two CNe that are considered as dubious by Livio and Truran (1994) on the basis of the high measured Ne abundances but relatively low values of the total heavy elements enrichment.  It should be noted that there were initially three dubious cases identified by Livio and Truran (1994), but IUE observations of Nova LMC\\,1990\\,N.1 (Starrfield et al.\\ 1992, Vanlandingham et al. 1999) confirmed it to be an ONeMg nova similar to V463\\,CrA.  \\begin{figure}[t] \\centering \\includegraphics[width=7cm]{f3.eps} \\caption{The histogram of the CNe distribution as a function of [Ne/O] abundance. The white area represents the CO novae, the shaded are represents the ONeMg novae, while the black area represents the two CNe which show extreme Ne enrichment but relatively small values for the total heavy elements enrichment. See text for more details.} \\label{histo} \\end{figure}   The value of [Ne/O]$>$1 determined in this paper for U\\,Sco places the binary among the ``true Ne novae'', hosting an ONeMg WD. It should also be noted that U\\,Sco was observed by IUE during the 1979 outburst and on that occasion Williams et al. (1981) reported absorption lines and P-Cyg profiles from C{\\sc iv}, Si{\\sc iv} and N{\\sc v} in their early epoch data. All ONeMg novae - contrary to the CO novae- show a P-Cyg profile phase in their UV spectra after the iron curtain phase and before the transition to the nebular spectrum (Shore 2008). The UV lines displayed at this stage are Si{\\sc iv} 1400, C{\\sc iv} 155 0, Al{\\sc iii} 1860, and Mg{\\sc ii} 2800 and show a saturated absorption trough with very high terminal velocity (Shore 2008). The IUE spectra of U\\,Sco taken +4 and +6 days after the maximum clearly show P-Cyg profiles with broad absorption troughs that are very  similar to the IUE spectra of the ONeMg novae discussed by Shore (2008). Hence, U\\,Sco should be regarded as a recurrent nova hosting a massive ONeMg WD (M$_{WD}\\sim$1.37-1.55 M$\\odot$, Hachisu et al.\\ 2000 and Thoroughgood et al.\\ 2001). The WD will undergo core collapse and will not explode as a SN\\,Ia, unless the current models about accreting ONeMg WDs are significantly in error and the U\\,Sco mass accretion rate and WD mass prove to be substantially smaller.  It is interesting to note that theoretical studies (e.g. Hachisu et al.\\ 2000; Livio 2000; Thoroughgood et al. 2002; see also Justham and Podsiadlowski 2008; Walder et al. 2010) have always looked at U\\,Sco and all RNe as likely progenitors of type Ia supernovae. However, observational works on this class of objects seem to show the opposite, i.e. that recurrent novae are not viable SN\\,Ia progenitors. On one hand, Della~Valle and Livio (1996) have shown that the frequency of RNe in the Milky Way, M31 and the LMC is significantly smaller (by $\\sim$1-2 orders of magnitude) than the supernova Ia rate deduced for these same galaxies. On the other hand, Selvelli et al.\\ (2008) have provided observational evidence that T\\,Pyx ejects more material than it accretes and therefore it cannot explode as a SN\\,Ia. This paper concludes that U\\,Sco, hosting a massive ONeMg WD, cannot explode as a SN\\,Ia either. Therefore, among RNe, only symbiotic recurrent novae ``survive'' as the possible progenitors of SN\\,Ia  (Justham and Podsiadlowski 2008; Di Stefano 2010), explaining at least some of the observed SN\\,Ia (Patat et al.\\ 2011; Di Stefano 2010). Still, the ultimate fate of RS\\,Oph, the prototype object of the symbiotic RNe, remains uncertain (e.g. Osborne et al.\\ 2006; Justham and Podsiadlowski 2008).  While the question of how many different stellar systems produce type Ia supernovae  remains unsolved, it has been proven critical to establish not only the mass of the WD and the ejecta, but also the primary star composition in candidate SN\\,Ia progenitors.  In the case of RNe, in particular, it will be important to establish whether an ONeMg WD is peculiar to U\\,Sco or rather common to the RNe of the same type or to all recurrent novae, as suggested by  Webbink (1990)."
    },
    "1107/1107.3775_arXiv.txt": {
        "abstract": "We study the influence of the propagation effects on the mean profiles of radio pulsars using the method of the wave propagation in the inhomogeneous media describing by Kravtsov \\& Orlov (1990). This approach allows us firstly to include into consideration the transition from geometrical optics to vacuum propagation, the cyclotron absorption, and the wave refraction simultaneously. In addition, non-dipole magnetic field configuration, drift motion of plasma particles, and their realistic energy distribution  are taken into account. It is confirmed that for ordinary pulsars (period \\mbox{$P \\sim 1$ s}, surface magnetic field \\mbox{$B_0 \\sim 10^{12}$ G}) and typical plasma generation near magnetic poles (the multiplicity parameter $\\lambda  = n_{\\rm e}/n_{\\rm GJ} \\sim 10^3$) the polarization is formed inside the light cylinder at the distance $r_{\\rm esc} \\sim 1000 R$ from the neutron star, the circular polarization being \\mbox{$5$--$20$\\%} which is just observed. The one-to-one correspondence between the signs of circular polarization and position angle ($p.a.$) derivative along the profile for both ordinary and extraordinary waves is predicted. Using numerical integration we now can model the mean profiles of radio pulsars. It is shown that the standard $S$-shape form of the $p.a.$ swing can be realized for small enough multiplicity $\\lambda$ and large enough bulk Lorentz factor $\\gamma$ only. It is also shown that the value of $p.a.$ maximum derivative, that is often used for determination the angle between magnetic dipole and rotation axis, depends on the plasma parameters {and could differ from the rotation vector model (RVM) prediction.} ",
        "introduction": "\\label{aba:sec1} More than forty years after discovery, our understanding of pulsar phenomenon leaves an ambiguous impression. On the one hand, the key properties were understood almost immediately (see, e.g., the monographs by Manchester \\& Taylor 1977; Lyne \\& Graham-Smith 1998): the stable pattern of radio emission is related to the neutron star rotation, the energy source is the kinetic energy of its rotation, and the release mechanism has the electromagnetic nature. The mean profiles of radio pulsars are well described by the hollow cone model (see below). Within four last decades enormous {amount} of observational data concerning polarization and other morphological properties of mean profiles was collected (Rankin 1983, 1990; Johnston et al. 2007; Weltevrede \\& Johnston 2008; Hankins \\& Rankin 2010; Keith et al. 2010; Yan et al. 2011). However, there is no agreement on the mechanism of coherent radio emission (Beskin 1999; Usov 2006; Lyubarsky 2008).  Any self-consistent theory of pulsar radio emission must include at least three main elements. First, it should describe a plasma instability that produces coherent radio emission. Second, the saturation of this instability that determines the intensity of the outgoing radio emission. Third, once the radiation is produced, its polarization properties are modified by the interaction with magnetospheric plasma and fields. These propagation effects should be accounted for to make a quantitative comparison of the theoretical predictions for the generation of radio emission with observational data.  There are several proposed mechanisms for the initial instability: an unstable flow of relativistic electron-positron plasma flowing along curved magnetic field lines (Goldreich \\& Keeley 1971; Blandford 1975; Asseo et al. 1980; Beskin, Gurevich \\& Istomin 1993, hereafter BGI); an instability caused by boundedness of the region of open field lines (Luo, Melrose \\& Machabeli 1994; Asseo 1995); an instability connected with kinetic effects, that can be caused by non-equilibrium of the particles energy distribution function (mainly, anomalous Doppler effect in the region of cyclotron resonance); two-stream instability (Kazbegi, Machabeli \\& Melikidze 1991); the instability connected with the nonstationarity of plasma particle production in the region of its generation (Lyubarskii 1996).  The saturation mechanism, whose investigation requires involving the effects of a nonlinear wave interaction, is the most complex from the theoretical point of view. Therefore, it is not surprising that only a few researchers have managed to consider this question consistently (see, e.g., Istomin 1988). Finally, the processes of the wave propagation in pulsar magnetosphere have not yet been investigated with sufficient detail either, although the part of theory that includes the propagation processes can be constructed using the standard linear methods of plasma physics.  {Recall that the present interpretation of the mean profiles of radio pulsars is based on the ground of so-called hollow cone model (see, e.g., Manchester \\& Taylor 1977). Within this approach the directivity pattern is assumed to repeat the profile of the number density of secondary particles outflowing along the open field lines. As the secondary plasma cannot be generated near the magnetic pole (where the curvature photons radiating by primary beam propagate almost along magnetic field lines), the particle number density is to have the 'hole' in its space distribution (Sturrock 1971; Ruderman \\& Sutherland 1975).}  There are four assumptions in the hollow cone model: first, the emission is generated in the inner magnetospheric regions (where the magnetic field may be considered as a dipole); second, the emission propagates along the straight line; third, the cyclotron absorption may be neglected; and fourth, the polarization is determined at the emission point. Such basic characteristics of the received radio emission allow to determine the change of the position angle ($p.a.$) of the linear polarization along the mean profile (Radhakrishnan \\& Cocke 1969) \\begin{equation} p.a. = {\\rm arctan} \\left(\\frac{\\sin \\alpha \\sin \\phi} {\\sin \\alpha \\cos \\zeta \\cos \\phi - \\sin \\zeta \\cos \\alpha}\\right). \\label{p.a.} \\end{equation} Here $\\alpha$ is the inclination angle of the magnetic dipole to the rotation axis, $\\zeta$ is the angle between the rotation axis and the observer's direction, and $\\phi$ is the pulse phase.  As a result, the radiation beam width $W_{\\rm r}$ itself and its statistical dependence on the period $P$ can be qualitatively explained under these assumptions (Rankin 1983, 1990). {As the pulsar radio emission is highly polarized (Lyne \\& Graham-Smith 1998), one could check the validity of the relation (\\ref{p.a.}) as well. As is well-known, in many cases the observed $p.a.$ swing is in good agreement with this theoretical prediction. Besides, the polarization observations show that the pulsar radio emission consists of two orthogonal modes, i.e., of two components which $p.a.$ differ by $90^{\\circ}$ (Taylor \\& Stinebring 1986). It is logically to connect two such components with two normal modes, ordinary and extraordinary ones, propagating in a magnetoactive plasma (Ginzburg 1961).} It is not surprising that the hollow cone model in its simplest realization is currently widely used for quantitative determination of the parameters of neutron stars.  At the same time, it is well known that, in general, three main assumptions are incorrect. First of all, after the paper by Barnard \\& Arons (1986), it became clear that the ordinary wave {(i.e., the wave which electric field belongs to the plane containing the wave vector ${\\bf k}$ and the external magnetic field ${\\bf B}$)} does not propagate in a straight line, but deflects away from the magnetic axis. Subsequently, this effect was studied in detail by Petrova \\& Lyubarskii (1998, 2000), it was an important element in the BGI theory. The correction to relation (\\ref{p.a.}) connected with the aberration was determined by Blaskiewicz et al. (1991), but it was {rarely} used in analysis of the observational data as well.  Further, the cyclotron absorption that must take place near the light cylinder (Mikhailovsky et al. 1982) turns out to be so large that it will not allow the radio emission to escape the pulsar's magnetosphere (see, e.g., Fussel et al. 2003). Finally, the limiting polarization effect had not been discussed seriously over many years, although it was qualitatively clear that this effect must be decisive for explaining of high degree of circular polarization, typically $5$--$20$\\%.  Indeed, in the region of the radio emission generation located at $10$--$100$ neutron star radii (these values result from the hollow cone model), the magnetic field is still strong enough so the polarization of the two orthogonal modes is indistinguishable from a linear one. {For this reason, it is logical to conclude that the polarization characteristics cannot be formed precisely in the emission region, and the propagation effects are to play important role (cf. Mitra et al. 2009). Nevertheless}, in an overwhelming majority of the papers, Eqn. (\\ref{p.a.}) is used to investigate the polarization.  Recall that the limiting polarization effect is related to the escape of radio emission from a region of dense plasma, where the propagation is well described in the geometrical optics approximation (in this case, the polarization ellipse is defined by the orientation of the external magnetic field in the picture plane), into the region of rarefied plasma, where the emission polarization becomes almost constant along the ray. This process was well studied (Zheleznyakov 1977; Kravtsov \\& Orlov 1990) and was used successfully for numerous objects, for example, in connection with the problems of solar radio emission (Zheleznyakov 1964). However, in the theory of pulsar radio emission, such problem has not been solved. Above the papers where the level $r = r_{\\rm esc}$ at which the transition from the geometrical optics approximation to the vacuum occurs, was only estimated (see, e.g., Cheng \\& Ruderman 1979; Barnard 1986), one can note only a few paper by Petrova \\& Lyubarskii (2000) (these authors considered the problem in the infinite magnetic field), by Petrova (2001, 2003, 2006), as well as the recent papers by Wang, Lai \\& Han (2010, 2011).  The goal of our paper is to consider all three main effects (i.e., refraction, cyclotron absorption, and limiting polarization) simultaneously in a consistent manner for realistic case. Not only the plasma density but also the magnetic field decreases with increasing distance from the neutron star will be included into consideration. Also, the non-dipole magnetic field, the drift motion of plasma particles, and realistic distribution function of outgoing plasma will be taken into account.  In section 2 both ordinary and extraordinary waves propagation in the pulsar magnetosphere is briefly considered. In addition, the hydrodynamic derivation of dielectric tensor of relativistic magnetized plasma is given. In section 3 the main parameters of our model are discussed. In section 4 we discuss the transition from geometrical optics to vacuum propagation, the cyclotron absorption, and the wave refraction. The results of numerical calculations of the mean profiles of radio pulsar are presented in section 5. Finally, in section 6 we discuss our main results.  It is necessary to stress that the main goal of this paper is in describing the theoretical ground of propagation effects only. For this reason, we are not going to discuss here in detail the theoretical predictions for real objects. This will be done in the separate paper. On the other hand, the theory described below is independent on the emission mechanism and, hence, can be applied to any theory of the pulsar radio emission. ",
        "conclusions": "In this paper we study the influence of the propagation effects on the mean profiles of radio pulsars. The Kravtsov-Orlov approach allows us firstly to include into consideration the transition from geometrical optics to vacuum propagation, the cyclotron absorption, and the wave refraction simultaneously. Arbitrary non-dipole magnetic field configuration, drift motion of plasma particles, and their realistic energy distribution were taken into account. Using numerical integration, {we found how the propagation effects can correct the main characteristics of the hollow cone model.}  \\begin{figure*} \\includegraphics[scale=0.32]{CoreConal.pdf} \\caption{Formation of the directivity pattern in the pulsar magnetosphere. The X-mode propagating rectilinearly forms the core component, and deflected from the magnetic axis O-mode forms the conal one. The lower curves on each top panels indicate the Stokes parameter $V$, and the bottom panels shows the $p.a.$ swing. Corresponding pulsar profiles was taken from Weltevrede \\& Johnston (2008)} \\label{CCmod} \\end{figure*}  To summarize, one can formulate the following results: \\begin{itemize} \\item We confirm the one-to-one correlation between the signs of circular polarization $V$ and the position angle derivative (${\\rm d} p.a./{\\rm d} \\phi$) along the profile for both ordinary and extraordinary waves. For the X-mode the signs should be the SAME, and OPPOSITE for the O-mode. \\item The standard $S$-shape form of the ${\\it p.a.}$ swing (1) can be realized for small enough multiplicity $\\lambda$ and large enough bulk Lorentz-factor $\\gamma$ only. In other cases the significant differences can take place. The location of the $p.a.$ maximum derivative $({\\rm d}p.a./{\\rm d}\\phi)_{\\rm max}$ is shifted to the trailing side of the pulse. \\item The value of $p.a.$ maximum derivative $({\\rm d}p.a./{\\rm d}\\phi)_{\\rm max}$, that is often used for determination the angle between magnetic dipole and rotation axis, depends on the plasma parameters [and differs from rotation vector model (RVM) value] and, hence, cannot be used without more precise specifying of magnetospheric plasma model. \\item In general, the trailing side of the emission beam is absorbed. \\end{itemize}  {In our opinion, even the preliminary consideration demonstrates that these conclusions are not in contradiction with observational data. The first statement is in good agreement with observations (Han et al. 1998;  Andrianov \\& Beskin 2010). In some cases, the sign reversal in the core can occur (see Fig.~\\ref{core} and Wang et al. 2010), that is detected for several pulsars (see also Radhakrishnan \\& Rankin 1990). The second point is in agreement with empirical model of Blaskiewicz et al. (1991). On the other hand, as the shift value depends on plasma parameters (assuming fixed geometry), the solution of the inverse problem (i.e., the fitting of the real pulsar profiles) can in principle provide us a chance of correct enough estimating of plasma parameters. Finally, the last point is in agreement with observations as well (see, e.g., Backus et el., 2010). In more detail, it takes place for large enough multiplicity parameter $\\lambda > 10^{3}$. But in some cases the leading part can be absorbed as well. It happens when the polarization forms close to the light cylinder.}  {Finally, Fig.~\\ref{CCmod} demonstrates our understanding of the X- and O-modes propagation in the pulsar magnetosphere. Here the dot-dash lines going to the panels with pulsar profiles show different intersections of directivity pattern, that form the observable profiles. Resulting from the different propagation, the O- and X-modes produce two concentric cones, the inner (core) part corresponding to the X-mode and the outer (conal) one to the O-mode. Remember that the analysis of the relative intensity in O- and X-modes is a question for generation theory. In the context of the present propagation theory this parameter should be considered as a free one.}  {Another source of uncertainty which can affect the formation of the mean profile (and, in particular, results in the formation of its more complicated shape) is the possible ''spotty'' of the directivity pattern, i.e., the presence of the separate radiative domains within the diagram (Rankin 1993). As is well-known, such separate domains are observed as a subpulse drift (Ruderman \\& Sutherland 1975, Deshpande \\& Rankin 2001).}  {In addition, the full analysis is to include the possible linear depolarization. E.g., depolarization at the edges of the profile can be easily interpreted as a mixing of X and O-modes (Rankin \\& Ramachandran 2003). But the total degree of linear polarization (e.g., if we clearly see only one mode) is a question to propagation in the generation region. If this region is rather thick then one can expect the possible depolarization. For this reason we do not treat here the total degree of linear polarization at all.}  As we see, in general if only one mode dominates in pulsar profile, then we expect statistically $D$ (double) mean profile for the O-mode and $S$ (single) one for the X-mode. Of course, we suppose here that there is nonzero natural width of the radiative domains that efficiently smooths the hollow cone corresponding to the narrow X-mode. Nevertheless, as is shown on Fig.~\\ref{CCmod}, the X-mode can in principal produce $D$ profile and O-mode can produce $S$ ones. {The $X_{D}$ and $X_{S}$ profiles can be obtained when only X-mode is produced in the formation region}. On the other hand, if there are two modes with the similar intensity, the triple (T) and multiple (M) profiles can be obtained as well. So, this picture is in agreement with Rankin's morphological classification (Rankin 1983). Our investigation provides an additional information to morphological structure, i.e., the information of the type of the mode that forms mainly the mean profile. As was demonstrated, only the correlation between $p.a.$ derivative and $V$ signs provides us this chance.  {Finally, the determination of the angular dimension of the core and conal beams depends greatly on the emission radius (i.e., again, on the generation mechanism). So the detailed analysis is also beyond the scope of current investigation (one can find empirical results on this in, e.g., Maciesiak \\& Gil 2011).  Thus, the approach developed in this paper allows the predictions of the theory of radio emission to be quantitatively compared with the observational data. Moreover, based on the observational data, one can reject a number of models in which the shapes of the profiles of the position angle and the degree of circular polarization never realized in practice are obtained. Thus, it becomes possible to solve the inverse problem, i.e., to determine the parameters of the outflowing plasma and the magnetic field structure.  We will be glad to any collaboration with observers for investigating the observational features of the mean profiles on the base of our theory. The source code of our calculations (in Mathematica 7 or C) is available under request."
    },
    "1107/1107.2918_arXiv.txt": {
        "abstract": "We analyze the stellar mass-to-light ratio (\\ML) gradients in a large sample of local galaxies taken from the Sloan Digital Sky Survey, spanning a wide range of stellar masses and morphological types. As suggested by the well known relationship between mass-to-light (\\ML) ratios and colors, we show that \\ML\\ gradients are strongly correlated with colour gradients, which we trace to the effects of age variations. Stellar \\ML\\ gradients generally follow patterns of variation with stellar mass and galaxy type that were previous found for colour and metallicty gradients. In late-type galaxies \\ML\\ gradients are negative, steepening with increasing mass. In early-type galaxies \\ML\\ gradients are shallower while presenting a two-fold trend: they decrease with mass up to a characteristic mass of $\\mst \\sim 10^{10.3} \\Msun$ and increase at larger masses. We compare our findings with other analyses and discuss some implications for galaxy formation and for dark matter estimates. ",
        "introduction": "\\label{sec:intro}  Colour and absorption line gradients are markers of stellar population variations within galaxies, providing important clues to galaxy evolution (e.g. \\citealt{Spolaor+09}; \\citealt{GCK11}; \\citealt{Rawle+10}; \\citealt{Spolaor+10}; \\citealt[herafter T+10]{Tortora+10CG}, \\citealt{Tortora+11CGsim}). Although spectral features are less ambiguous probes of stellar population properties than colours are (e.g. \\citealt{Spolaor+09,Spolaor+10}; \\citealt{Kuntschner+10}; \\citealt{Rawle+10}), there are only small samples available outside the nuclear regions of galaxies. Fortunately, studies based on colours turn out to provide, on average, similar age and metallicity estimates to spectral studies (e.g. T+10), allowing for much larger surveys of radially-extended stellar populations.   In T+10 we analyzed colour gradients within an effective radius (\\Re) for an extensive sample of low-redshift galaxies based on SDSS data (\\citealt{Blanton05}), including both early-type (ETGs) and late-type (LTGs) galaxies. We carried out simple stellar populations analyses, fitting synthetic models to the observed colours to derive estimated age and metallicity gradients, and stellar masses (see T+10 and later in the paper for further details). The main results found can be summarized as follows: \\begin{enumerate} \\item Colour gradients are, on average, negative (i.e. galaxies are redder in the centre), except for galaxies with lower stellar masses  ($\\mst \\lsim 10^{9.3} \\, \\rm \\Msun$) which tend to have null or positive gradients. At fixed mass, LTGs have steeper gradients than ETGs, whose colour profiles in general show milder variations . \\item The colour gradients of LTGs decrease systematically with mass while the trend for ETGs changes near a mass of $\\mst \\sim 10^{10.3} \\, \\rm \\Msun$: systematically decreasing and increasing at lower and higher masses, respectively. \\item The primary driver of the colour gradients, mainly for galaxies with older centers, is metallicity variations, with age gradients playing a secondary role. For LTGs, the metallicity gradients tend to be steeper for higher masses, The same is true for ETGs except that for masses above $\\mst \\sim 10^{10.3} \\, \\rm \\Msun$, the (negative) gradients are fairly invariant or slightly increasing with mass. \\item At fixed mass, galaxies with older centres have the shallowest metallicity and age gradients. \\end{enumerate}  \\begin{figure} \\psfig{file=fig1.eps, width=0.48\\textwidth} \\caption{Diagram of colour versus stellar mass for the local SDSS galaxy sample analyzed in T+10. The grey-scale contours show the density of data points, with brighter regions marking higher densities. The thin white line separates the red sequence and blue cloud. The arrows and the curves in the inset panel give information about the efficiency of the phenomena which drive the two fold trend we discuss in the paper (see T+10 for further details). AGN feedback and merging are important at high mass with an efficiency that increases with mass, while SN feedback and gas inflow drive the galaxy evolution at lower masses, with an efficiency that is larger at the lowest masses.} \\label{fig:fig1} \\end{figure}  We found that these results were generally consistent with others in the literature (e.g. \\citealt{Spolaor+09}; \\citealt{Rawle+10}) once different galaxy selection effects were taken into account.  The overall observational picture can be framed within the predictions of hydrodynamical and chemodynamical simulations of galaxy formation (\\citealt{BS99}; \\citealt{Kawata01}; \\citealt{KG03}; \\citealt{Ko04}; \\citealt{Hopkins+09a}; \\citealt{Pipino+10}; \\citealt{Tortora+11CGsim}) which include various physical processes that can influence the gradients: supernovae and AGN feedback, galaxy mergers, interactions with environment, etc. A fresh and qualitative overview of this physical scenario is given by Fig. \\ref{fig:fig1}, where we show the isodensity contours for the colour-mass diagram (which we have already discussed in T+10) and superimpose some lines (in the inset) and arrows which indicate the efficiency of the highlighted physical processes. In particular, the observations for high-mass ETGs seem to generally suggest the influence of mergers and AGN feedback (\\citealt{Ko04}; \\citealt{Hopkins+09a}; \\citealt{Sijacki+07}), with an efficiency possibly increasing with mass (see red arrows and lines in the inset panel). On the other hand, the population gradients of LTGs and less massive ETGs may be mainly driven by infall and supernovae feedback (\\citealt{Larson74, Larson75}; \\citealt{Kawata01}; \\citealt{KG03}; \\citealt{Ko04}; \\citealt{Pipino+10}; \\citealt{Tortora+11CGsim}), which are more efficient at very low masses (see blue lines and arrows).  In this paper we consider the implications of the observed colour, metallicity and age gradients from T+10 for variations of stellar {\\it mass-to-light ratio} $M/L$ with radius.  This issue is relevant for understanding the detailed distribution of mass within galaxies:  both the stellar mass and the dark matter, whose distribution is inferred by substracting the stellar mass from the total dynamical mass. There has been scarcely any literature to date that focussed on this topic, except for a study in spiral galaxies by \\cite{PS10}.   As a first approximation, one could use single-colour semi-empirical models such as those of \\cite{BdJ01} and \\citet[B03 hereafter]{Bell+03} to go directly from colour gradients to \\ML\\ gradients. However, we wish to make use of multiple colours, to allow for departures of individual galaxies from mean trends, to explore the systematics thoroughly, and to trace the dependence on stellar parameters, like age and metallicity.    This paper is organized as follows. The sample and the spectral analysis are discussed in \\S \\ref{sec:sample}, in where we also show the ability of our fitting procedure to recover the correct \\ML\\ gradients, and we examine possible systematics. The results are shown in \\S \\ref{sec:results}, while \\S \\ref{sec:conclusions} is devoted to a discussion of the physical implications and conclusions. ",
        "conclusions": "\\label{sec:conclusions}     In this paper we have investigated the correlation between \\ML\\ gradients and stellar mass or velocity dispersion for a large sample of local galaxies from SDSS. The gradients \\gML\\ have been derived by the fitting of synthetic spectral models to observed optical colours. We have used Monte Carlo simulations to check that the optical-only wavelength baseline should not dramatically affect the results compared to including near-infrared and ultraviolet data. This conclusion is encouraging in view of future studies of higher redshift galaxies which are generally observed in rest-frame optical bands. Anyway, future analysis in local samples, relying on near-IR and/or UV data would be able to give further details about the degeneracies and systematics affecting stellar fitting.  We have found that there exists a tight positive correlation between \\gML\\ and colour gradients which is not surprising because of the well known \\ML\\--colour correlations (e.g., \\citealt{BdJ01}; B03). We have found that \\ML\\ is a decreasing function of stellar mass in LTGs, while a two-fold trend is found for ETGs (for both stellar mass and velocity dispersion), in agreement with the trends of colour gradients discussed in T+10. Similar results are obtained using the \\ML\\--colour relations from B03, which, although easy to use, do not give information about stellar population parameters. Thus, the novelty of our approach resides in our ability to check how \\gML\\ gradients are driven by variations in metallicity or age. In general, we have found that age is most important, but that metallicity becomes more important in many high-mass galaxies.   This picture does not change dramatically if \\ML\\ gradients are calculated in redder bands ($V$- and $i$-bands rather than $B$). In these cases \\gML\\ are shallower and the trends with stellar mass and velocity dispersion become milder (Fig. \\ref{fig:fig3}). These results suggest that \\ML\\ estimates from red or infrared bands would not depend on galactocentric radius, since the radiation emitted by stellar populations in these spectral regions is more homogenous.  This is the first work to derive \\ML\\ gradients 1) from observations of a large galaxy sample and 2) using a direct stellar population analysis approach. The only point of comparison is the work by \\cite{PS10} in a sample of spiral galaxies, where negative \\ML\\ gradients were found to be important in modelling rotation curves.   We may explain some of our observed trends through physical mechanisms as follow. Starting from low-mass systems, centrally-young LTGs have age increasing with radius while metallicity decreases, which can be due to the presence of a gas inflow which allows the formation of younger stars in the centre with a smaller \\ML. Centrally-old LTGs have age decreasing with radius which implies recent outer star formation with a similar metallicity to the centre. Here the central star formation might have been quenched by SNe which have expelled metals to larger radii (where the younger stars drive a smaller \\ML). Centrally-young ETGs have positive age, metallicity and \\ML\\ gradients, and may hosting expanding shells (\\citealt{Mori+97}). Centrally-old ETGs show mild negative age gradients and very shallow positive metallicity gradients, similarly to the old LTGs thus again suggesting SNe having played a role.  The more massive centrally-young LTGs have decreasing \\ML\\ with radius, but also strong negative \\gZ, thus the new stars formed inside the core (\\gage$>0$) have higher metallicity and \\ML. This might be due to recycled gas from SNe or AGN which is falls back to the centre due to the deep potential wells and produces a rejuvenated stellar population. In centrally-young ETGs, on the other hand, the younger stars formed later in the centres (i.e. with positive age gradients) and with higher metallicities, possibly because of a central burst of star formation in ``wet'' mergers (e.g. \\citealt{RJ04}; \\citealt{Hopkins+09a}; \\citealt{Hopkins+10}; T+10). Centrally-old LTGs and ETGs show very similar behaviors, which seem compatible with merging events that largely wash out any age gradients (and the \\gZ\\ in ETGs), and may produce coeval metallicities in the outer regions of the LTGs. In both cases the net effect is that metallicity is the main driver of the strong negative \\gML\\ in LTGs and of the weak negative \\gML\\ of the massive ETGs (in any). Thus, the physical mechanisms that can account for the observations outlined in Tab. \\ref{tab:tab2} seem consistent with the evolutionary scenario discussed in \\S\\ref{sec:intro}.   In particular, the two-fold trend of ETGs in Fig. \\ref{fig:fig3}, with a minimum around $\\log\\mst/\\Msun=10.3$ is mirrored by the colour and stellar population gradients as in Fig. \\ref{fig:fig1} (and T+10) and found also in other works (\\citealt{Rawle+10}; \\citealt{Spolaor+09, Spolaor+10}). The characteristic mass agrees well with the typical mass scale break for star forming and passive systems (\\citealt{Kauffmann2003}), for ``bright'' and ``ordinary'' galaxies (\\citealt{Capaccioli92a}; \\citealt{GG03}; \\citealt{Graham+03}; \\citealt{Trujillo+04}; \\citealt{Kormendy+09}), or for cold and hot flow (shock heating; \\citealt{dek_birn06}).  Following the main conclusions in T+10, quasi-monolithic collapse and supernovae feedback could be the main drivers of the decreasing \\gML-\\mst\\ trends in LTGs and low mass ETGs (\\citealt{Larson74}; \\citealt{Kawata01}, \\citealt{KG03}), while the increasing trends in massive ETGs could be explained by the heightened importance of mergers and AGN feedback (\\citealt{Ko04}; \\citealt{Sijacki+07}; \\citealt{Hopkins+09a}).   \\ML\\ gradients are expected from simulations of galaxy formation and evolution. E.g. \\cite{BP99} found in chemo-dynamical simulations of spiral galaxies that the central regions ($< 12 \\, \\rm kpc$) have steep negative \\ML\\ gradients, becoming flatter at larger radii. High-redshift ($z\\sim 2$) merger remnants simulated by \\cite{Wuyts+10} showed time-dependent negative \\ML\\ as a consequence of redder cores in redder galaxies (similarly to observations of local galaxies in T+10). \\ML\\ gradients were also predicted for merger remnants by \\cite{Hopkins+10}, who stressed that the effective radius derived from the $B$-band light would increase with time (by $\\sim$~10\\%) simply because of stellar population gradients. This effect is important for conclusions about stellar densities and related dynamical implications (e.g. for the fundamental plane). Similarly \\citet{van Dokkum08} argued that the existence of such gradients can imply that the observed quiescent galaxies at large redshift may be of an order or magnitude more dense than previously found from observations.  The \\ML\\ gradients present a general complication for estimates of the dark matter distribution in galaxies. \\citet{PS10} found that for spiral galaxies, including \\ML\\ gradients would change the fits to circular velocity curves but not produce strong effects on the inferred dark matter density parameters. To our knowledge, this issue has not been addressed in the literature for early-type galaxies, and it is beyond the scope of the present paper to do so in any detail. In particular, we have studied the stellar population gradients only inside 1~\\Re, while the gradients out to $\\sim$~5~\\Re\\ could be important for dark matter analyses (e.g. \\citealt{2005MNRAS.357..691N, Napolitano+09, Napolitano+11}). Assuming that the gradients do not change radically outside $\\sim$~1~\\Re, for the massive ETGs generally studied with lensing or dynamics ($\\sig_0 \\gsim$~150~km~s$^{-1}$), we would expect the stellar \\ML\\ gradients to be mild (see Fig.~\\ref{fig:fig3}), and to not strongly affect the dark matter estimates.  A final issue is that until now, we have assumed that the IMF is invariant. If there are galaxy-to-galaxy variations in the IMF (see e.g. \\citealt{2010Natur.468..940V,vDC11,2011arXiv1104.2379G}), then the stellar \\ML\\ values are changed only by a constant multiplicative value, and the implied gradients remain unchanged. However, IMF variations if real would then likely exist {\\it within} galaxies owing to variations in stellar populations. Indeed, there have been many suggestions that stars forming at early times had a non-standard IMF (e.g. \\citealt{2005MNRAS.359..211L,2007MNRAS.374L..29K,van Dokkum08,2008MNRAS.385..147D,Holden10}).  One study of nearby ETGs suggested that systems with older stars within 1~\\Re\\ have ``lighter'' (closer to Chabrier than Salpeter) IMFs \\citep{NRT10}. This would imply that positive and negative age gradients would decrease and increase the stellar \\ML\\ gradients, respectively. Revisiting Fig. \\ref{fig:fig3}, the gradients for the centrally-young objects would decrease and the centrally old objects would increase, possibly reducing and even eliminating the age-related scatter in the gradient trends.  Another scenario would be if the central regions of massive ETGs have ``higher-mass'' IMFs than standard (\\citealt{vDC11}), while the outer regions have accreted galaxies with more ``normal'' IMFs (e.g. \\citealt{Kroupa01} or \\citealt{Chabrier03}). This would decrease these galaxies' stellar \\ML\\ gradients, although it is not clear how this effect would relate to age and metallicity trends.   In future analyses, we plan to analyze stellar \\ML\\ profiles out to a few \\Re, and analyze the impact of \\ML\\ gradients on dark matter profile inferences, which may be compared to $\\Lambda$CDM (e.g. \\citealt{NFW96}). We will also consider the (probably small) adjustments implied for central dark matter inferences (e.g. \\citealt{Tortora10lensing}). Dust effects will further be examined: dust is not important for metallicity gradients (see T+10) but could affect age and thereby \\ML\\ gradients, particularly for intermediate-mass ETGs."
    },
    "1107/1107.2129_arXiv.txt": {
        "abstract": "{We report on the detection of bright CO(4-3) line emission in two powerful, obscured quasars discovered in the SWIRE survey, SW022513 and SW022550 at z$\\ge$3.4. We analyze the line strength and profile to determine the gas mass, dynamical mass and the gas dynamics for both galaxies. In SW022513 we may have found the first evidence for a molecular, AGN-driven wind in the early Universe. The line profile in SW022513 is broad (FWHM\\,=\\,1000\\,\\kms) and blueshifted by $-$200\\,\\kms\\ relative to systemic (where the systemic velocity is estimated from the narrow components of ionized gas lines, as is commonly done for AGN at low and high redshifts). SW022550 has a more regular, double-peaked profile, which is marginally spatially resolved in our data, consistent with either a merger or an extended disk.  The molecular gas masses, $4\\times 10^{10}$\\,\\msun, are large and account for $<$30\\% of the stellar mass, making these obscured QSOs as gas rich as other powerful CO emitting galaxies at high redshift, i.e., submillimeter galaxies. Our sources exhibit relatively lower star-formation efficiencies compared to other dusty, powerful starburst galaxies at high redshift. We speculate that this could be a consequence of the AGN perturbing the molecular gas.} ",
        "introduction": "The most accredited galaxy evolution models postulate that the progenitors of massive galaxies evolve through a phase of intense star-formation activity at high redshift, accompanied by accretion onto a central supermassive black hole~\\citep[e.g.][]{sanders88,granato01,granato04,dimatteo05}. Star formation and accretion are both fueled by cold gas, and are therefore intricately linked. Several cosmological models now postulate that the phase of growth of black hole and host is terminated when the accreting black hole is able to heat and expel the ambient gas ~\\citep[e.g.,][]{sanders88,granato01,granato04,dimatteo05,hopkins05b,churazov05,merloni08,fanidakis11}. Given the central role of cold gas in these processes, molecular emission line studies should be ideal to test and refine these scenarios~\\citep[e.g.,][]{narayanan09}.  Owing to their brightness, CO emission lines are commonly used to trace the immense reservoirs of molecular gas in high-redshift galaxies, providing valuable constraints on the gas content, gas heating and cooling processes, and on the evolutionary and dynamical state of these systems. CO line profiles, and in particular the molecular gas kinematics observed with interferometers at high spatial resolution, allow to differentiate between, e.g., rotating disks and mergers. Under the assumption that the gas is driven by gravity and is approximately virialized, CO line profiles and velocity gradients can be used to estimate dynamical masses. In recent years, CO observations have led to a plethora of new constraints, e.g., on star-formation efficiencies in different types of high-$z$ galaxies~\\citep{greve05,solomon05,tacconi08,coppin08b,iono09,daddi10b,riechers11}, the regulation of star-formation in the most rapidly growing galaxies~\\citep{beelen04,greve05,daddi08,schinnerer08,nesvadba09a,genzel10}, or the co-evolution of black holes and their host galaxies~\\citep{peng06b,maiolino07,alexander08,riechers09a,wang10}.  Molecular gas studies have been carried out in starburst-dominated galaxies, e.g., submillimeter galaxies (SMGs), type I QSOs, radio galaxies, and even normal star-forming galaxies over large redshift ranges~\\citep[e.g.,][]{omont96,neri03,solomon05,weiss05,greve05,carilli06,tacconi08,coppin08b,nesvadba09a,riechers09a,riechers09b,daddi10b,bothwell10,riechers11}. By comparison, only few obscured (type II) quasars have been targeted in CO millimeter studies~\\citep{aravena08,martinez09,yan10}. This is unfortunate, because obscured quasars could represent the short, but critical transition phase between starburst and type I (unobscured) QSO activity~\\citep[e.g.,][]{sanders88,hopkins06b}, but has also been unavoidable so far, because identifying these objects and obtaining spectroscopic redshifts remains a challenge.  The {\\it Spitzer Space Telescope}~\\citep{werner04} unveiled a significant number of obscured quasars owing to their bright mid-infrared (MIR) emission. However, only about 20-25\\% of these galaxies are also bright at millimeter wavelengths~\\citep{lutz05,polletta09,martinez09}, and only few of those have accurate spectroscopic redshifts, which enable detailed studies of their millimeter CO emission~\\citep{sajina08,yan10}. Those that have been detected in CO have large molecular gas masses and a wide range of CO line profiles~\\citep{aravena08,martinez09,yan10}. Their overall properties are similar to SMGs, suggesting they may be mergers~\\citep{dasyra08,yan10} with intense, dusty star formation. Obscured quasars are expected to be close to the `blowout' phase when the AGN heats and expels the remaining cold ISM, terminating star formation and accretion activity. However, no signature of such winds has so far been found, although broad CO line wings have now been identified in a growing number of nearby AGN, illustrating that molecular outflows are indeed possible~\\citep{feruglio10,alatalo11}.  Here we present CO millimeter observations of two obscured QSOs associated with massive, intensely star-forming galaxies. These systems are amongst the most luminous obscured QSOs at high-$z$ currently known. They are thus ideal candidates to investigate if obscured QSOs are indeed close to complete their phase of active growth, and in particular, if this phase is terminated by feedback from their AGN. We will argue that this could potentially be the case in one of our targets. Throughout the paper, we adopt a $\\Lambda_{CDM}$ cosmology with H$_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$=0.73, and $\\Omega_{M}$=0.37~\\citep{spergel03}.  \\begin{table}[h!] \\caption{\\label{co_params}Redshifts, masses, luminosities, and CO parameters.} \\centering \\begin{tabular}{lrr} \\hline\\hline Parameter&SW022550&SW022513\\\\ \\hline $z_{UV}$                                         &    3.867$\\pm$0.009 &            \\nodata \\\\ $z_{opt}$                                        &    3.876$\\pm$0.001 &    3.427$\\pm$0.002 \\\\ F$_{1.2mm}$ [mJy]                                &      4.70$\\pm$0.77 &      5.53$\\pm$0.72 \\\\ M$_{stellar}$\\tablefootmark{a} [10$^{10}$\\msun]  &             16--80 &             20--40 \\\\ Log(L(IR)\\tablefootmark{a}) [\\lsun]              &         12.5--13.3 &         12.5--13.2 \\\\ SFR\\tablefootmark{a} [\\msun/yr]                  &          500--3000 &          500--3000 \\\\ $z_{CO}$                                         &  3.8719$\\pm$0.0008 &  3.4220$\\pm$0.0007 \\\\ $z_{CO,1}$                                       &  3.8686$\\pm$0.0007 &  3.4208$\\pm$0.0013 \\\\ $z_{CO,2}$                                       &  3.8761$\\pm$0.0012 &  3.4270$\\pm$0.0007 \\\\ $\\nu$\\tablefootmark{b}  [GHz]                    &            94.6330 &           104.2598 \\\\ $\\nu_1$\\tablefootmark{b} [GHz]                   &            94.6978 &           104.2885 \\\\ $\\nu_2$\\tablefootmark{b} [GHz]                   &            94.5504 &           104.1432 \\\\ FWHM\\tablefootmark{b} [\\kms]                     &        800$\\pm$120 &       1020$\\pm$110 \\\\ FWHM$_1$\\tablefootmark{b} [\\kms]                 &         290$\\pm$90 &        950$\\pm$150 \\\\ FWHM$_2$\\tablefootmark{b} [\\kms]                 &        420$\\pm$180 &        200$\\pm$150 \\\\ $\\Delta$v\\tablefootmark{c} [\\kms]                &         460$\\pm$90 &        420$\\pm$100 \\\\ I$_{CO}$ [Jy\\,\\kms]                              &      1.46$\\pm$0.28 &      2.62$\\pm$0.36 \\\\ I$_{CO,1}$ [Jy\\,\\kms]\\tablefootmark{d}           &      0.68$\\pm$0.11 &      2.25$\\pm$0.28 \\\\ I$_{CO,2}$ [Jy\\,\\kms]\\tablefootmark{d}           &      0.70$\\pm$0.23 &      0.28$\\pm$0.21 \\\\ L$_{CO}$\\tablefootmark{e} [10$^{8}$\\,\\lsun]      &      1.77$\\pm$0.32 &      2.61$\\pm$0.36 \\\\ L$^{\\prime}_{CO}$\\tablefootmark{f} [10$^{10}$\\,K\\,\\kms\\,pc$^2$]& 5.7$\\pm$1.0 & 8.3$\\pm$1.2 \\\\ M$_{gas}$\\tablefootmark{g} [10$^{10}$\\,\\msun]    &     4.53$\\pm$0.82  &      6.66$\\pm$0.92 \\\\ M$_{gas}$/M$_{stellar}$                          &        0.06--0.29  &        0.17--0.33  \\\\ M$_{dyn}^{disk}\\times sin^2(i)$ [10$^{10}$\\,\\msun]  &      40$\\pm$20  &        2--34       \\\\ M$_{dyn}^{merger}\\times sin^2(i)$ [10$^{10}$\\,\\msun]&      20$\\pm$10  &        9--35       \\\\ SFE\\tablefootmark{h} [\\lsun/(K\\,\\kms\\,pc$^2$)]   &            35--194 &           24--108  \\\\ \\hline \\end{tabular} \\tablefoot{All CO parameters refer to the $J$=(4--3) transition.\\\\ \\tablefoottext{a}{L(IR) is the 8--1000\\,$\\mu$m luminosity in \\lsun, M$_{stellar}$ the stellar mass, and SFR the star-formation rate~\\citep[see~\\S~\\ref{src_desc} and][]{polletta08c}.}\\\\ \\tablefoottext{b}{$\\nu$ is the observed frequency of the CO line fitted with a single Gaussian. $\\nu_1$ and $\\nu_2$ are the observed frequencies of the two CO peaks obtained with a double Gaussian fit, and FWHM, FWHM$_1$, and FWHM$_2$ are the corresponding full-widths-at-half maximum.}\\\\ \\tablefoottext{c}{Velocity offset between the peaks of the double Gaussian fit.}\\\\ \\tablefoottext{d}{Intensity derived from each component of the double Gaussian fit.}\\\\ \\tablefoottext{e}{L$_{CO}$ is the CO line luminosity, in\\,\\lsun.}\\\\ \\tablefoottext{f}{L$^{\\prime}_{CO}$ is the integrated CO luminosity, in K\\,\\kms\\,pc$^2$.}\\\\ \\tablefoottext{g}{M$_{gas}$ is the gas (H$_2$+He) mass.}\\\\ \\tablefoottext{h}{SFE is the star-formation efficiency defined as L(FIR)/L$^{\\prime}_{CO}$ (see~\\S~\\ref{sfe}).} } \\end{table} ",
        "conclusions": "We presented the properties of the molecular gas and the possible role of powerful AGN activity on this gas in two highly luminous obscured QSOs at $z$$>$3.4.  Our study was based on observations of the CO(4--3) line carried out with the PdBI. The luminosity and profile of the CO line indicate that both systems contain large amounts of molecular gas (M$_{gas}$$\\gtrsim$4$\\times$10$^{10}$\\,\\msun) corresponding to $<$30\\% of their stellar mass. The CO lines are among the broadest CO lines ever observed ($\\sim$800--1000\\,\\kms). A double-peaked line is observed in SW022550, and a blueshifted broad profile in SW022513. Various scenarios were discussed to explain the observed line profiles and gas kinematics: a rotating disk, a merger, and an AGN-driven outflow.  The CO properties in SW022550 indicate that the molecular gas might be distributed in a highly inclined, almost edge-on, rotating disk or be associated with a merger. The CO emission appears extended on scales of $\\sim$16\\,kpc, consistent with both scenarios. The disk inclination is constrained by the presence of two peaks in the line profile and by the high estimates of rotational velocities and dynamical masses. A merger can also explain the CO line kinematics, even if there is no evidence of multiple sources in the available optical and infrared images. A comparison with the rest-frame optical spectrum of the source~\\citep{nesvadba11b} suggests that the molecular gas is not significantly affected by the presence of the AGN.  For SW022513, the disk and merger hypotheses seem less convincing. Based on an analysis of the rest-frame optical line emission in SW022513, we suspect that parts of the molecular gas could be stirred up by the AGN and perhaps be part of an AGN-driven outflow. This is further supported by a detailed comparison with NGC1266, the clearest example of an AGN-driven molecular wind in a nearby galaxy. In this scenario, most of the line flux would be produced in a very turbulent and potentially outflowing wind component. The wind hypothesis would very naturally explain the similarity of the CO and H$\\beta$ line profiles.  A comparison of our objects' molecular gas properties with those of other powerful high-redshifts systems, i.e., submillimeter galaxies, type I QSOs, RGs, and $z$$\\gtrsim$2 ULIRGs, shows that they are characterized by relatively lower star-formation efficiencies, defined as the L(FIR)/L$^{\\prime}_{CO}$ luminosity ratio. Similarly low SFEs are also found in other obscured AGNs from the literature, but they are not typical of this class of objects. Another difference between our sources and SMGs and type I QSOs it is that they already seem to lie near the local BH-bulge mass relation. These results suggest that our sources, and perhaps other obscured QSOs, might be near the end of their most active growing phase and on the verge of turning off and transitioning into passively evolving massive galaxies. Their low star-formation efficiencies might be due to the perturbation caused by the AGN on the molecular gas.  Our observations indicate that the AGN plays a major role in perturbing the molecular gas in SW022513. This source might thus provide a laboratory to test and constrain evolution models and to study the formation of massive galaxies at critical early epochs. It could also be used as reference to search for similar sources and carry out more targeted searches for molecular gas outflows."
    },
    "1107/1107.4069_arXiv.txt": {
        "abstract": "Waves reflected by the inner edge of a protoplanetary disk are shown to significantly modify Type I migration, even allowing the trapping of planets near the inner disk edge for small planets in a range of disk parameters. This may inform the distribution planets close to their central stars, as observed recently by the Kepler mission. ",
        "introduction": "The location and configuration of planets within their planetary systems depend sensitively on the effects  of many physical processes including planet-planet scattering \\citep[e.g.][]{Rasio96}, Kozai oscillations \\citep[e.g.][]{Fabrycky07}, and disk migration \\citep[e.g.][hereafter GT80]{GT80}\\footnote{For review and comparison of migration mechanisms see \\citet{Morton11} and \\citet{Ward00}, and references therein.}. Of these, disk migration is likely to have been dominant in systems where planets are observed to be co-planar \\citep[e.g.][]{Steffen10}. The density waves launched inwards from the inner Lindblad resonances carry away negative angular momentum driving the planet outwards (positive torque). The waves launched outwards from the outer Lindblad resonances carry away positive angular momentum driving it inwards  (negative torque). For planets too small to clear a gap in their protoplanetary disk  Type I migration occurs due to the asymmetry between the Lindblad torques \\citep[hereafter W86 and W97]{Ward86, Ward97} with the outer torque being typically stronger, driving a net migration inwards.  Recent observations by the Kepler spacecraft have shed new light on to the distribution of extrasolar planets within their planetary systems. \\citet{Howard11} examined the radial distribution of planets closer than 0.25 AU of solar type stars observed in the Kepler field. They found that smaller planets ($R_p \\leq 2-4 R_{\\earth}$) have an exponential cutoff in the radial distribution at an orbital period $ \\sim 7$ days, while larger planets have exponential cutoffs closer to the star at $\\sim 2$ days.  The protoplanetary disks are likely be truncated before direct contact with the host stars by a strong stellar magnetic field  near to the corotation radius \\citep[see e.g.][]{Shu94}. Rotational studies of classical T Tauri stars give periods ranging from $\\sim 3-7$ days for stars with $M_* > 0.25 M_\\odot$ \\citep{Lamm05}. The truncation radius is typically taken to be $\\sim 5$ stellar radii, or about $ \\sim 4$ day periods for typical 2 solar radius pre-main-sequence stars \\citep{Gullbring98}. The inner edge of the accretion disk plays a role in halting migration very close to the host star; disk migration is naturally halted once a planet passes through the inner disk edge \\citep{Lin96} and the Lindblad resonances can no longer interact with the disk material.  Close to the inner edge of the disk the effect of the corotation resonance also becomes important and can exceed the differential Lindblad torque, particularly at locations with strong gradients in the surface density profile \\citep{GT79}. However, unless the viscosity is strong enough to transfer angular momentum away from the co-orbital region quickly enough, the corotation resonance saturates and the corotation torque goes to zero \\citep{Masset01, Ogilvie03, Baruteau08}.  In this letter we will outline a halting mechanism for Type I migration of planets due to density waves resonant between the inner Lindblad resonance and the disk inner edge. This allows for the Type I migration of smaller planets to be halted within the disk, at a few times the inner disk radius. ",
        "conclusions": "We have shown that for small planets the linear response of a disk resonance caused by an inner reflecting edge can act to trap planets at a few times the inner disk radius. The exact location of this trapping depends on the details of the disk parameters and dissipation mechanism. For larger planets, the disk will likely undergo a non-linear resonance, which typically has non-linear resonant gain decreasing with increasing planet mass. This implies that above a certain cutoff mass, the resonant gain due to the non-linear resonance will be insufficient to halt migration, and the planets will continue to migrate inwards through the inner disk edge. There they will be halted at $r_{\\rm halt} $ when the lowest order outer Lindblad resonance is located within the inner edge of the disk. The corotation torque is also important near these edges, but is susceptible to saturation, unless the viscosity is sufficiently high. If unsaturated the corotation torque is mainly active when the planet is very close to the disk edge and the density gradient (and thus vortensity gradient) are strongest. However, further from the disk edge where the wave trapping resonance is active, both effects may be important in determining the net torque.  This naturally generates two populations of planets, with halting radius a few times the inner disk edge for smaller planets, and with larger planets halted inside the inner disk edge. This process may help to explain the observations of planets close to solar type stars by \\citet{Howard11} which showed that larger planets were halted with typical periods of $\\sim 2$ days, corresponding to $r_{\\rm halt} \\sim 0.03$ AU, while smaller planets were halted with typical periods of $\\sim 7$ days, $r_{\\rm halt} \\sim 0.07$ AU, consistent with inner disk edges with periods of  $\\sim 4$ days corresponding to $r_{\\rm in} \\sim 0.05$ AU.  This process may also be important for planets located near other reflecting edges within the disk, such as the edges of magnetic dead zones \\citep[see e.g.][]{Gammie96}, or the edge of a gap in the disk opened by a giant planet \\citep[see e.g.][]{Lin79}.  The effect of turbulent density fluctuations on this wave resonance may be to ratchet the system through successive ``traps'', if the variation is on a scale similar to the separation between the Lindblad resonances, such that the relative amplitudes of the Lindblad torques are affected. Larger resonant peaks will be more difficult to overcome, however a detailed study of the effect of turbulence is beyond the scope of the current work."
    },
    "1107/1107.5223.txt": {
        "abstract": "It is only now, with low-frequency radio telescopes, long exposures with high-resolution X-ray satellites and $\\gamma$-ray telescopes, that we are beginning to learn about the physics in the periphery of galaxy clusters. In the coming years, Sunyaev-Zeldovich telescopes are going to deliver further great insights into the plasma physics of these special regions in the Universe. The last years have already shown tremendous progress with detections of shocks, estimates of magnetic field strengths and constraints on the particle acceleration efficiency. X-ray observations have revealed shock fronts in cluster outskirts which have allowed inferences about the microphysical structure of shocks fronts in such extreme environments. The best indications for magnetic fields and relativistic particles in cluster outskirts come from observations of so-called radio relics, which are megaparsec-sized regions of radio emission from the edges of galaxy clusters. As these are difficult to detect due to their low surface brightness, only few of these objects are known. But they have provided unprecedented evidence for the acceleration of relativistic particles at shock fronts and the existence of \\muG\\, strength fields as far out as the virial radius of clusters. In this review we summarise the observational and theoretical state of our knowledge of magnetic fields, relativistic particles and shocks in cluster outskirts. ",
        "introduction": "\\includegraphics[width=1\\textwidth]{entr.eps} \\caption{Normalized post-shock gas entropy, $K_{\\rm mf} = K_0\\, T/\\rho^{(\\gamma_{\\rm g} -1)}$, as a function of the CR particle energy escape flux $Q_{\\rm esc}$ carried by energetic particles. The dimensionless flux is defined as $\\eta_{\\rm esc}= Q_{\\rm esc}/\\Pi_{\\rm kin}$, where $\\Pi_{\\rm kin} = \\rho_1 v_{\\rm sh}^3/2$. The upper curve (dotted) corresponds to an effective adiabatic exponent $\\gamma_{\\rm g}$ = 4/3 (relativistic gas), while the lower (solid) curve corresponds to $\\gamma_{\\rm g}$ = 5/3. The postshock entropy is normalized to $K_{\\rm mf}(\\eta_{\\rm esc}=0)$.} \\end{figure*}  In galaxy clusters, entropy profiles have been measured by \\citet{prattea10} at least out to $R_{1000}$ in 31 nearby galaxy clusters from the \\emph{XMM-Newton} cluster structure survey (REXCESS) and out to $R_{500}$ in thirteen systems. The effects of CRs on the entropy production in shocks illustrated above can be of importance for  accretion shocks that are expected to be located at larger radii. It will require deeper exposures of the next generation telescopes to measure the entropy profiles at these radii. The entropy in the inner cluster regions can be also affected by internal shocks if the shocks are efficient accelerators of CRs. The current entropy profile simulations \\citep[e.g.][]{borganiea08,bk09} ought to be extended to include CR physics. ",
        "conclusions": "It is only now, with low-frequency radio telescopes, long exposures with high-resolution X-ray satellites and $\\gamma$-ray telescopes, that we are beginning to learn about the physics of cluster outskirts. In the coming years SZ telescopes are going to deliver further insights into the plasma physics of these special regions in the Universe. The last years have already shown tremendous progress with detections of shocks, estimates of magnetic field strengths and constraints on the particle acceleration efficiency. The main points of this review are listed below:  \\begin{itemize}  \\item Cosmological shocks are collisionless: Coulomb collisions are not sufficient to provide the viscous dissipation of the incoming flow, and collective effects play a major role.  \\item Collisionless shocks generate energetic, non-thermal particles that can penetrate far into the upstream flow. The particles decelerate the flow and preheat the gas. They can also efficiently generate strong fluctuating magnetic fields in the upstream region. This turbulence, which may result from cosmic-ray driven instabilities, in the shock precursor, is certain to play a critical role in non-linear models of strong, CR-modified shocks.  \\item A few merger shocks have been identified in X-rays, exhibiting both a sharp gas density edge and an unambiguous temperature jump: in the \"bullet cluster\", 1E 0657\u00d056 \\citep{markevitch02}, A520 \\citep{markevitch05} and Abell 2146 \\citep{russell10}.  \\item On the basis of X-ray observations, microphysical properties of shocks can be derived. Markevitch et al. (2006) find that the temperatures across merger shocks are consistent with instant heating; equilibration between electrons and protons on the collisional timescale is excluded. The electron temperature downstream of the shock depends on collisionless electron heating e.g. \\citep{bu99,rlg08,bpp08}.  Non-resonant interactions of the electrons with strong nonlinear fluctuations generated by kinetic instabilities of the ions in the transition region inside the shock front may play the main role in the heating and pre-acceleration of the electrons. Future X-ray observations extending to the virial radius or even close to the shock radius should be able to detect signatures of non-equipartition at shocks.   \\item \\emph{Suzaku} observations provide evidence for departures from hydrostatic equilibrium around and even before the virial radius as it is seen in the galaxy cluster Abell 1795 \\citep[\\emph{e.g.,} ][]{bautzea09}.  Evidence for nonthermal pressure support suggests that bulk motions from mergers could be making a significant contribution to the gas energy in the outskirts of this cluster \\citep{georgeea09}.  \\item Recently discovered radio relics provide very strong evidence for shock acceleration at merger shock waves in galaxy clusters. Van Weeren et al. (2010) find that on average smaller radio relics have steeper spectra. Such a trend is in line with predictions from shock statistics derived from cosmological simulations \\citep{skillman08, battaglia09,hoeft08}. They find that larger shock waves occur mainly in lower-density regions and have larger Mach numbers, and consequently shallower spectra.  \\item Observations of radio relics in combination with X-ray observations suggest magnetic field strengths in relics of $>3$ \\muG\\, and efficiencies of 0.2\\% of the kinetic energy dissipated in a shock going into relativistic electrons (Finoguenov et al. 2009, van Weeren et al. 2010) .  \\item Radio relics also suggest that the efficiency of shock acceleration at weak shocks is higher than predicted in theories of diffusive shock acceleration.  \\item The bulk of the energy in their simulations is dissipated in galaxy clusters which contribute about 75 \\% of the total energy dissipation (about 80 \\% the contribution from cluster outskirts is included), while filaments contribute about 15 \\% of the total energy dissipation. In agreement with observational findings, the maximum diffuse radio emission of galaxy clusters depends strongly on their X-ray temperature. While shocks around $M \\sim 2$ contribute most to the generation of gas thermal energy, shocks around $M \\sim 3$ contribute most to the generation of CR energy \\citep{ryu03,krco07}.  \\item The effects of CRs on entropy production in shocks as illustrated above can be of importance for accretion shocks that are expected to be located at larger radii. It will require deeper exposures of the next generation telescopes to measure the entropy profiles at these radii.  \\end{itemize} Many open questions remain. So far still little is known about shock acceleration and magnetic field generation in clusters, nor about the microstructure of shocks in very dilute plasmas.  Most likely, all these processes are related and their understanding requires input from, both, observations and theory. The next years will see a rapid growth in observational data. On the theoretical side, we expect substantial progress from new techniques that simulate particle acceleration in a magnetohydrodynamical setting. Fully kinetic particle-in-cell simulations provide a powerful tool for exploration of the structure of collisionless shocks from first principles, thus determining self-consistently the interplay between shock-generated waves and accelerated particles \\citep{sironi10}. Results from these simulations can then serve as input for cosmological simulations that include CR-physics."
    },
    "1107/1107.0967_arXiv.txt": {
        "abstract": "Multi-wavelength observations/spectroscopy of exoplanetary atmospheres are the basis of the emerging exciting field of comparative exoplanetology. The \\hr planetary system is an ideal laboratory to study our current knowledge gap between massive field brown dwarfs and the cold 5-Gyr old Solar system planets. The \\hr planets have so far been imaged at J- to L-band, with only upper limits available at M-band. We present here deep high-contrast Keck II adaptive optics M-band observations that show the imaging detection of~$3$ of the~$4$ currently known \\hr planets. Such detections were made possible due to the development of an innovative LOCI-based background subtraction scheme that is~$3$ times more efficient than a classical median background subtraction for Keck II AO data, representing a gain in telescope time of up to a factor of~$9$. These M-band detections extend the broad band photometric coverage out to~$\\sim 5\\mu$m and provide access to the strong CO fundamental absorption band at~$4.5\\mu$m. The new M-band photometry shows that the \\hr planets are located near the L/T-type dwarf transition, similar to what was found by other studies. We also confirm that the best atmospheric fits are consistent with low surface gravity, dusty and non-equilibrium CO/CH$_4$ chemistry models. ",
        "introduction": "After more than a decade of searching, the direct exoplanet imaging quest was finally successful in 2008 with the discovery of three planetary systems \\citep{marois08,kalas08,lagrange09}. One key advantage of this technique is the detection of the planet's thermal emission. Detailed multi-band photometry and spectrometry can be acquired that can then be compared with atmospheric models to derive the planet's physical characteristics and study the effect of dust and molecular chemistry. With its multiple co-eval Jovian planets, the \\hr system is an ideal laboratory to study young planets with low temperature/surface gravity atmospheres and bridge the gap between massive field brown dwarfs and the cold Solar system planets.  \\hr is a 30Myr old~\\citep{marois10,zuckerman11} A5V star located 39.4pc away \\citep{leeuwen07} in the Pegasus constellation. It is classified as a~$\\gamma$ Doradus and a $\\lambda$ Bootis star~\\citep{gray99}. It also shows an IR excess~(Vega-like star) consistent with a debris disk composed of a warm dust disk~(6 to 15AU), a massive cold dust disk~(90 up to 300AU) and a small dust particle halo extending up to a 1,000AU \\citep{rhee07,su09}.  Photometry and some spectroscopy of the \\hr planets are available in several bands from 1 to 3.8$\\mu$m~\\citep{marois08, lafreniere09,metchev09,janson10,marois10, hinz10,currie11,barman11}. An initial characterization \\citep{marois08,bowler10,marois10,currie11,barman11} has been achieved using the available data and state-of-the-art atmospheric models developed for field brown dwarf analysis. It is clear that these planets are quite different than most field brown dwarfs, showing very dusty atmospheres while being cool ($\\sim $1,000K) with no sign of methane molecular absorption.  L- and M-band detections/upper limits suggest this lack of methane is due to a CO/CH$_4$ non-equilibrium chemistry \\citep{saumon03,cushing06,leggett07,hinz10,bowler10,currie11,barman11}. Other teams have been able to fit the available photometry with patchy clouds and equilibrium chemistry atmospheres ~ \\citep{marois08,currie11}. Accurate M-band photometry can help disentangle the various models, but it is hard on a ground-based telescope due to the bright thermal background mainly from the telescope and the adaptive optics system~\\citep{Lloyd00}. This background is particularly hard to remove at Keck, as it varies with time as the instrument image rotator (which is located very near focus) moves and as the AO deformable mirror modulates the rotator-induced background pattern as seen by the detector. As a result, conventional sky/background subtraction routines leaves a spatially variable residual that limits final sensitivity.  In this letter, we present the first M-band ($4.670 \\mu$m) detections of three of the currently four known \\hr planets, the longest wavelength at which these planets have been imaged. The observations are discussed in~\\S\\ref{sec : obs}. The data reduction along with the new very efficient LOCI-based (LOCI: locally optimized combination of images) background subtraction technique are presented in~\\S\\ref{sec : reduc}. Comparisons with field brown dwarfs and new atmospheric fits are shown in \\S\\ref{sec : dis}. Conclusions are summarized in~\\S\\ref{sec : conc}. ",
        "conclusions": "\\label{sec : conc} In this letter, we have estimated for the first time the M-band fluxes of three of the currently four known \\hr planets. These detections were made possible due to the use of an innovative LOCI-based background subtraction routine that has allowed for a factor of~$3$ gain in contrast (factor of~$9$ in integration time) compared to a classical background subtraction using a median. This new background subtraction routine can be used to subtract the background noise in any infrared data.  We have detected \\hrbnosp, c and~d at M-band from 3 to 8$\\sigma$. From a K'-L' and L'-M' color diagram, we confirm that the three planets are located near the end of the L-type sequence, close to the L- and T-dwarf transition region. We then derived new atmosphere model fits for the three planets. For planets~c and~d, temperatures and surface gravities are close to the expected values from evolutionary models. For planet~b, the solar abundance model fits well the broad band photometry from~$1$ to~$5\\,\\mu$m but it requires a very small planetary radius to match the bolometric luminosity and is thus in contradiction with the planet evolution models. The metal rich evolution-consistent model over-predicts the M-band flux, which may indicate an even higher overall metallicity and/or a non-solar C-to-O ratio. Higher~SNR images will help disentangle the different physical and chemical parameters."
    },
    "1107/1107.2570.txt": {
        "abstract": "Star formation is arguably the most important physical process in the cosmos. It is a fundamental driver of galaxy evolution and the ultimate source of most of the energy emitted by galaxies.  A correct interpretation of star formation rate (SFR) measures is therefore essential to our understanding of galaxy formation and evolution.  Unfortunately, however, no single SFR estimator is universally available or even applicable in all circumstances: the numerous galaxies found in deep surveys are often too faint (or too distant) to yield significant detections with most standard SFR measures, and until now there have been no global, multi-band observations of nearby galaxies that span all the conditions under which star-formation is taking place.  To address this need in a systematic way, we have undertaken a multi-band survey of all types of star-forming galaxies in the local Universe.  This project, the Star Formation Reference Survey (SFRS), is based on a statistically valid sample of 369 nearby galaxies that span all existing combinations of dust temperature, SFR, and specific SFR.  Furthermore, because the SFRS is blind with respect to AGN fraction and environment it serves as a means to assess the influence of these factors on SFR.  Our panchromatic global flux measurements (including \\GG\\ FUV+NUV, SDSS $ugriz$, 2MASS $JHK_s$, \\SS\\ 3--8\\,$\\mu$m, and others) furnish uniform SFR measures and the context in which their reliability can be assessed.  This paper describes the SFRS survey strategy, defines the sample, and presents the multi-band photometry collected to date. ",
        "introduction": "Star formation has been the most important single physical process since recombination.  Not only have stars created most of the luminous energy in the Universe, they have produced the heavy elements needed for planets and life.  Star formation has naturally been the subject of vast numbers of studies, leading, for example, to our current understanding of the star formation history of the Universe (\\eg, Lilly et al. 1996, Hopkins \\& Beacom 2006; Madau et al. 1998) and the discovery of the Schmidt Law (Schmidt 1959; updated by Kennicutt et al. 1998) relating the SFR in galaxies to the gas surface density.  A general inadequacy of existing work is the lack of a self-consistent treatment of SFR across the full electromagnetic spectrum: ultraviolet (UV) continuum only samples the SFR in the absence of dust; \\halpha\\ and [\\ion{O}{2}] measure {\\em ionizing} photons coming from only the high-mass ($\\ga$5\\,\\Msun) population.  Emission in the polycyclic aromatic hydrocarbon (PAH) bands is taken as a SFR measure (\\eg, Madden \\etal\\ 2006; Wu \\etal\\ 2005) but may be unreliable for low-luminosity (or low-metallicity) galaxies (Hogg \\etal 2005) or in the presence of a hard radiation field.  Far-infrared dust re-radiation samples a broad range of star formation ($\\sim$1--10\\,\\Msun\\ for reasonable IMFs) but only becomes a precise SFR measure in the optically-thick limit (Kennicutt 1998).  Nonthermal radio emission comes from Type II supernova remnants (\\eg, Rieke \\etal 1980) and therefore only samples higher-mass SF.  Star formation activity is by no means evenly distributed throughout galaxies.  Instead, stars form within the densest regions of giant molecular clouds; this phenomenon manifests in the \\IRAS\\ bands as far-infrared radiation when the (hot) young stellar objects deeply embedded in their natal clouds illuminate the surrounding interstellar material (\\eg, Parker 1991), which then reradiates that energy at long wavelengths.  In a detailed case study of Milky Way GMC complexes (\\eg, the California Nebula) Lada, Lombardi, \\& Alves (2010) demonstrated a remarkably tight correlation between SFR and the total mass of {\\sl dense} gas, lending further support to this view.  There is considerable evidence in the literature that our understanding of star formation depends on making consistent use of all available wavebands.  For example, Kartaltepe \\etal (2010) showed that without photometry at wavelengths longer than 100\\,$\\mu$m, total far-IR luminosities (and thus SFRs) are typically underestimated by 0.2\\,dex; in some cases the discrepancy can be much larger.  This can be interpreted as an inability to adquately measure emission from a cold dust component (if present) when such far-IR data are lacking. These issues are now becoming better understood as a result of \\Akari\\ and \\HH\\ programs that reach deeper into the far-IR than could \\IRAS, \\ISO, or \\SS. The situation remains complex, however. In \\HH-selected galaxies, dust attenuation appears to strongly impact the UV detection fraction and the relationship between UV and IR SFR indicators (Buat \\etal 2010). There are strong hints of systematic discrepancies between \\Akari\\ and \\IRAS\\ photometry at $\\sim$100\\,$\\mu$m (Jeong et al. 2007; Figs.~1 and 7 of Takeuchi et al 2010). These unresolved issues can in principle be addressed with thoughtful controls.  But there are also systematic effects intrinsic to the galaxies themselves.  For example, there is evidence that the specific SFR (sSFR, the SFR per unit stellar mass) depends on stellar mass (\\eg, Sobral \\etal 2011, Elbaz \\etal 2011).  For many years there have been suggestions that dust temperature was dependent on luminosity (\\eg, Sanders \\etal 2003).  And there is as yet no systematic treatment of the contribution of the older, quiescent stellar components to the SFR tracers most commonly used.  To better understand the complexities of star formation, a comprehensive treatment examining the influence of all the major parameters --- stellar mass, dust mass and temperature, and metallicity --- on SFR indicators is needed. The importance and timeliness of this subject is attested by numerous recent efforts to grapple with the nuances of SFR estimation. Notable examples include the \\SS\\ Infrared Nearby Galaxies Sample (SINGS; Kennicutt \\etal 2003; Calzetti \\etal 2010), the Local Volume Legacy Survey (LVLS; Dale \\etal 2009), the \\HH\\ Reference Survey (HRS; Boselli \\etal 2010), the Great Observatories All-sky LIRG Survey (GOALS; Armus \\etal 2009), and the Multi-Wavelength Extreme Starburst Sample (MESS; Laag \\etal 2010), among others. Each of these undertakings has been designed to attack specific aspects of the star-formation phenomenon, but each also has limitations that prevent it from providing a comprehensive picture of star formation in galaxies.  This paper presents the Star Formation Reference Survey (SFRS), a sample of nearby galaxies having a unique capacity to describe star formation under all conditions in which it occurs in the local Universe. Section~\\ref{sec:selection} describes the selection criteria used to define the SFRS.  Section~\\ref{sec:datasets} presents the SFRS datasets obtained to date.  Section~\\ref{sec:discussion} shows how the SFRS complements the above-mentioned projects, presents an estimate of AGN prevalence and the ramifications for far-infrared SFR measures, and briefly describes SFRS-related observing campaigns now in progress. ",
        "conclusions": "\\label{sec:conclusion}  The Star Formation Reference Survey, by virtue of its reliance on a restricted but representative far-IR selection, ensures the capability to study {\\sl obscured} star formation in all of its varied manifestations in the local Universe.  Its panchromatic resources, spanning UV to radio wavelengths and featuring 100\\% complete photometry in the visible to mid-infrared regimes ($ugrizJHK_s$ and four IRAC bands) provides an opportunity to quantitatively assess the degree to which far-IR emission reflects total (not just obscured) SFR.  This comprehensive collection of the most widely-used SFR indicators will furnish an invaluable resource for the interpretation of more distant galaxies, \\ie, galaxies for which some or even nearly all such SFR metrics are inaccessible because of their relative faintness, so that a context exists in which to better understand the limited data available. Because the SFRS is fully representative of star-forming galaxies in the local Universe, it is an optimal benchmark for understanding star formation in the distant cosmos.  The interplay of various star formation and AGN indicators will be explored in future papers.  Paper~2 will examine the relationship of UV SFR indicators to those obtained in the other bands. P. Bonfini \\etal (2012, in preparation) will describe the methods whereby relatively faint AGNs in the sample are identified using structural decomposition. Y.-N. Zhu \\etal (2012, in preparation) will examine a suite of narrow-band \\halpha\\ imaging and compare the global \\halpha-derived SFRs to those in the other bands compiled for the SFRS.  All the related photometry will be made public to facilitate investigations by others in the community."
    },
    "1107/1107.0821_arXiv.txt": {
        "abstract": "The Tornado nebula (G\\,357.7$-$0.1) is a mysterious radio source with bright ``head'' and faint ``tail'' located in the direction of the Galactic center (GC) region. We here report the discovery of two diffuse X-ray sources at the head and tail of the Tornado with the Suzaku satellite. We found emission lines from highly ionized atoms in the two sources. The spectra are reproduced by an optically thin thermal plasma with a common temperature of 0.6--0.7~keV. The interstellar absorption ($N_{\\rm H}$) of these sources are the same and are slightly larger than that of the GC distance. Since the estimated distance using the $N_{\\rm H}$ value is consistent with the radio observation of the Tornado, these X-ray sources are likely associated with the Tornado nebula. The twin-plasma morphology at the both ends of the Tornado suggests that the system is a bipolar/outflow source. ",
        "introduction": "\\label{sec:intro}  G\\,357.7$-$0.1 is a bright radio object near the Galactic center (GC) direction and was first catalogued by \\citet{1960AuJPh..13..676M} as MSH\\,17-3{\\it9}. The extended, non-thermal emission with the spectral index of $\\sim -0.6$ (\\cite{1973AuJPh..26..379D}, \\cite{1974IAUS...60..347S}) and the linear polarization of $\\lesssim$10\\% (e.g. \\cite{1974AJ.....79..132K}) placed this source as one of the Galactic supernova remnant (SNR).  It has an elongated shape with $\\sim$\\timeform{10'} length and has the brightest peak offset to one side, which is called the ``head'', while fainter emissions are elongated to the other side, called the ``tail''. Unusual structure was found with the higher resolution observations. The remarkable feature is  the  several co-axial filamentary arcs swirling around its major axis, and hence nicknamed as the ``Tornado nebula'' (figure~\\ref{fig:xisimg}a: \\cite{1985Natur.313..113S}, \\cite{1985Natur.313..115B}, \\cite{1985Natur.313..118H}). \\citet{1994ApJ...432L..39S} corrected the large-scale brightness gradient of the radio map, and found a bipolar radio structure along the major axis.  The distance to the Tornado is constrained to be $> 6$~kpc from H\\emissiontype{I} absorption measurements (\\cite{1972ApJS...24...49R}). Recently, OH maser at 1720 MHz, supporting evidence for  the shocked molecular clouds, is detected at the head (\\cite{1996AJ....111.1651F}, \\cite{1999ApJ...527..172Y}, \\cite{2004MNRAS.354..393L}).  From the velocity of the OH maser ($-12.4$~km~s$^{-1}$), the distance is estimated to be 11.8~kpc (\\cite{1996AJ....111.1651F}). The H$_2$ emission lines are also detected toward the head of the Tornado (\\cite{2004MNRAS.354..393L}).  The origin of the Tornado nebula has been controversial due to its peculiar appearance. Some scenarios are  based on the head-tail structure; the center of activity is located at the brightest head and the tail is the periphery structure. The source may be either a radio galaxy with one-sided jet (\\cite{1980A&A....90..269W}, \\cite{1980ARA&A..18..165M}), a Galactic nebula powered by a pulsar, or a Galactic binary source located at the head (\\cite{1985Natur.313..113S}, \\cite{1985Natur.313..115B}, \\cite{1989ApJ...346..860S}).  The other scenarios are along the lines of bipolar-like structure: an extragalactic double-lobed source (\\cite{1989PASAu...8..184C}), an exotic remnant by an equatorial supernova of a rotating massive star (\\cite{1985Natur.313..113S}), or precessing jets of binary accretion system (\\cite{1985Natur.313..113S}, \\cite{1987A&A...171..205M}, \\cite{1989PASAu...8..184C}). The problem of these scenarios is that no compact source has been found near at the center of the Tornado.  \\begin{table*} \\caption{Log of Suzaku observations of the Tornado}\\label{tab:obslog} \\begin{center} \\begin{tabular}{cccccc} \\hline \\hline Sequence no. & \\multicolumn{2}{c}{Aim point} & Start date & Effective & Field name \\\\ & $\\alpha$ (J2000.0) & $\\delta$ (J2000.0) & & exposure & \\\\ \\hline 503015010 & \\timeform{17h40m07s} & \\timeform{-30D57'45''} & 2008/9/19 & 56.8~ks & obs08 \\\\ 504036010 & \\timeform{17h40m31s} & \\timeform{-30D56'56''} & 2009/8/29 & 125~ks & obs09 \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*} \\begin{center} \\FigureFile(160mm,90mm){figure1.eps} \\end{center} \\caption{(a) The 1.4~GHz radio continuum map from the Very Large Array data archives. Overlaid contours represent the radio map with the Australia Telescope Compact Array (ATCA) made by combining 4.79 and 5.84~GHz images (the fainter emission along the tail is enhanced: \\cite{1994ApJ...432L..39S}). The XIS fields of view are indicated with the solid and dashed rectangles. (b) Smoothed broad-band XIS images of the Tornado nebula: soft (0.5--1.5~keV in red), medium (1.5--3.0~keV in green), and hard (3.0--7.0~keV in blue) bands are displayed. The data from the three CCDs are merged, NXB are subtracted, and vignetting are corrected. The two fields are combined with normalization of the exposure times. Source and background extraction regions are shown with the solid and dashed ellipses, respectively.}\\label{fig:xisimg} \\end{figure*}  In the X-ray band, ASCA found a hint of weak X-rays from the Tornado (\\cite{2003ApJ...585..319Y}). Then \\citet{2003ApJ...594L..35G} discovered a diffuse X-ray source in the head with Chandra. From the centrally filled X-rays, the authors proposed that the Tornado is a Galactic mixed-morphology SNR (\\cite{1998ApJ...503L.167R}). However the nature of the X-rays, whether thermal or non-thermal, was not clear due to the limited photon statistics.  In order to give further constraint on the origin from the high energy phenomena, we conducted deep X-ray observations on the Tornado nebula with the X-ray Imaging Spectrometer (XIS: \\cite{2007PASJ...59S..23K}) onboard Suzaku (\\cite{2007PASJ...59S...1M}). We present the discovery of twin thermal plasmas associated with the head and tail of the Tornado. We investigate possible interaction of the Tornado with molecular clouds by referring the data from the Nobeyama Radio Observatory (NRO) 45-m telescope. Based on these new results, we discuss the nature of the X-ray source.  Throughout this paper, east and west indicate  for the direction in the Galactic longitude, and north and south for that in the Galactic latitude. Statistical uncertainties are represented by the 90\\% confidence intervals. ",
        "conclusions": "\\subsection{Nature of the X-ray Sources and Association to the Tornado}\\label{sec:nature}  We find that the absorption column densities ($N_{\\rm H}$) for the NW and SE sources are very similar to each other (table 3).  Since the $N_{\\rm H}$ values for the GDXE near the source region are almost constant at $(5.3\\pm0.4)\\times 10^{22}$~cm$^{-2}$, the same $N_{\\rm H}$ values for NW and SE supports that the two sources are located at nearly the same distances.  Assuming that the distance of the GDXE is 8~kpc and that the interstellar gas density along the line-of-sight is constant, the distances of NW and SE can be estimated by the $N_{\\rm H}$ ratio relative to the background emission of ($5.3\\pm0.4)\\times 10^{22}$~cm$^{-2}$ as $10.0^{+1.8}_{-1.4}$ and $11.2^{+2.7}_{-2.3}$~kpc, respectively. These distances are almost consistent with  the radio distance of the Tornado ($\\sim$12~kpc; \\cite{1996AJ....111.1651F}). In the projected image, the radio emission of the Tornado appears to be associated with the two X-ray structures at the both ends (see figure~\\ref{fig:multiimg}). Thus we argue that the two X-ray sources are physically connected to the Tornado nebula.  The two diffuse sources are found to exhibit optically thin thermal nature. In fact, the X-ray emission from the sources are described with CIE plasmas of nearly the same temperatures (0.6--0.7~keV). The abundances are consistent with  the solar values. This indicates that the X-ray emitting materials mainly consist of interstellar gas. If the distance of the Tornado is 12~kpc, the X-ray source size of \\timeform{4'} corresponds to 14~pc.  Assuming a spherical shape with uniform density, we derive the physical parameters of the NW and SE plasmas, electron density ($n_{\\rm e}$), mass ($M$) and total thermal energy ($E_{\\rm th}$), as are listed in table~\\ref{tab:plasmapar}.  We found no hint of non-thermal emission from the NW and SE sources. Assuming the photon index of 2, we obtained the 3-$\\sigma$ upper limits of the luminosities of non-thermal X-rays to be $1.1$\\% and $2.0$\\% of those of the thermal plasmas in the 0.5--7.0~keV band for NW and SE, respectively.  \\begin{table}[!ht] \\begin{center} \\caption{Physical parameters of the X-ray emitting plasma.}\\label{tab:plasmapar} \\begin{tabular}{lcc} \\hline Parameter & NW & SE \\\\ \\hline $VEM$ ($10^{58}~$cm$^{-3}$)\\footnotemark[$*$]& 1.0 & 1.5 \\\\ $n_{\\rm e}$ (cm$^{-3}$) & 0.49 & 0.59 \\\\ $M$ ($M_{\\odot}$) & 23 & 29 \\\\ $E_{\\rm th}$ ($10^{49}$~erg) & 8.3 & 8.2 \\\\ \\hline \\multicolumn{3}{@{}l@{}}{\\hbox to 0pt{\\parbox{70mm}{\\footnotesize \\par\\noindent \\footnotemark[$*$] These values are taken from the best-fit model in table 3 at the distance of 12~kpc.}\\hss}} \\end{tabular} \\end{center} \\end{table}  \\begin{figure}[!h] \\begin{center} \\FigureFile(80mm,90mm){figure4.eps} \\end{center} \\caption{Multi-wavelength image of the Tornado nebula: the medium band (1.5--3.0~keV) image overlaid with contours of the radio continuum emission with the ATCA (blue) and CO $J$=1--0 ($V_{\\rm{LSR}}=-11\\pm1$~km~s$^{-1}$) with the NRO 45~m telescope (green). The FOV of the NRO observation is shown with the solid rectangle.} \\label{fig:multiimg} \\end{figure}  \\subsection{Bipolar Structure of the Tornado Nebula}\\label{sec:bipolar}  The two X-ray plasmas are located at the both ends (head and tail) of the symmetric (bipolar-like) radio structure found by \\citet{1994ApJ...432L..39S}. Since the properties of the X-ray plasmas are almost identical, these would be the twins formed by a common process. A plausible scenario is that the X-ray plasmas have been generated by the termination shocks  at the intersecting molecular clouds. In fact the shocked molecular cloud associated with the Tornado has been found  near NW (\\cite{1996AJ....111.1651F}).  To confirm the cloud interactions, we further search for the molecular structure near at the  NW and SE regions using the data of CO $J=$1--0 taken with the NRO 45-m telescope (private communication with T. Oka). The results in the velocity range of $11\\pm1$~km~s$^{-1}$, close to the velocity of the OH maser near at NW, are shown in figure~\\ref{fig:multiimg}. We clearly see CO counterparts near NW (MC\\,1) and SE (MC\\,2).  As we described in \\S~\\ref{sec:intro} (Introduction), many ideas for the origin of the Tornado have been proposed. The present result that the Tornado has bipolar X-rays at the head and tail excludes most of these scenarios. In the following sections (\\S~\\ref{sec:jet} and \\S~\\ref{sec:remnant}), we discuss two plausible origins.  \\subsection{Bipolar Jet from a Binary Accretion System}\\label{sec:jet}  One possibility is that the Tornado is made by bipolar jets from a binary accretion system. The filamentary radio arcs swirling around the major axis would be due to the precession of jets as is found in SS\\,433. The twin X-ray emitting plasmas are made by the shock with molecular clouds. Bipolar jets have usually a bright central power source, but no X-ray source is found from the center of the Tornado nebula. We set the 3-$\\sigma$ upper limit on the luminosity to be $1.9\\times 10^{33}$~erg~s$^{-1}$ in the 0.5--10.0~keV band, assuming the absorption column of $N_{\\rm H}=7\\times10^{22}$~cm$^{-2}$, the power-law spectrum with the photon index of 2, and the distance of 12~kpc. Chandra has gave more severe constraint of  $\\sim1\\times10^{33}$~erg~s$^{-1}$ (\\cite{2003ApJ...594L..35G}). These upper limits are only less than 1\\% of  NW and SE luminosities of $2.4\\times 10^{35}$~erg~s$^{-1}$ and $3.6\\times 10^{35}$~erg~s$^{-1}$, respectively. A plausible explanation for the non-detection of the bright central source is that putative binary went into a quiescent state after forming jet-structures in a past active phase.  The synchrotron cooling timescale (\\cite{1996ApJ...459L..13R}) for electrons which emit $\\epsilon_{\\rm syn}$~keV photons in a magnetic field $B$~G  is estimated to be, \\begin{eqnarray} \\tau \\approx 2000\\left( \\frac{B}{10\\ \\mu{\\rm G}} \\right)^{-\\frac{3}{2}} \\left(\\frac{\\epsilon_{\\rm syn}} {1\\ {\\rm keV}}\\right)^{-\\frac{1}{2}}\\ {\\rm years}. \\end{eqnarray} Assuming the magnetic field of $100$~$\\mu$G (\\cite{2008ApJS..177..255L}), the cooling timescale for $\\ge 1$ keV photons is estimated to be less than $\\sim$40--50~years. The lower energy electrons responsible for the radio synchrotron radiation can survive much longer time. The cooling time of the X-ray emitting plasma, which are made by the interaction between the jets and the molecular clouds, is also much longer. Thus the radio Tornado and the twin X-rays at the both ends would be fossil radiations of a past activity some  $\\sim$40--50~years ago. Although such object is very rare, we see a hint in the jets of 4U\\,1755$-$33. The central source is in low luminosity of $3.6\\times10^{31}$~erg s$^{-1}$ at present, but was bright two decades ago (\\cite{2003ApJ...586L..71A}, \\cite{1996IAUC.6302....2R}).  \\subsection{Remnant of an Equatorial Explosion}\\label{sec:remnant}  An alternative possibility for the bipolar structure of the Tornado is a remnant of an equatorial explosion of a rotating high mass star (\\cite{1985Natur.313..113S}). For example, a hypernova explosion associated with a $\\gamma$-ray burst may produce a bipolar blast wave (\\cite{2004ApJ...613L..17I}). If this blast wave hits molecular clouds, then the double X-ray spots will be produced. This scenario does not necessary require a central source. However, a hypernova is very energetic, and in our case, the explosion energy is less than $10^{51}$ erg s$^{-1}$ (table 4). We therefore suspect this possibility is less likely compared to the bipolar-jet model of an X-ray binary."
    },
    "1107/1107.2564.txt": {
        "abstract": "We have mapped the stellar and gaseous kinematics, as well as the emission-line flux distributions and ratios, from the inner $\\approx$\\,450\\,pc radius of the Seyfert 2 galaxy Mrk\\,1157, using two-dimensional (2D) near-IR $J-$ and $K_l-$band spectra obtained with the Gemini NIFS instrument at a spatial resolution of $\\approx$35~pc and velocity resolution of $\\approx$\\,40\\,\\kms.  The stellar velocity field shows a rotation pattern, with a discrete S-shaped zero velocity curve -- a signature of a nuclear bar. The presence of a bar is also supported by the residual map between the observed rotation field and a model of circular orbits in a Plummer potential. The stellar velocity dispersion ($\\sigma_*$) map presents a partial ring of low-$\\sigma_*$ values ($50-60$\\,km\\,s$^{-1}$) at 250\\,pc from the nucleus surrounded by higher $\\sigma_*$ values from the galaxy bulge. We propose that this ring has origin in kinematically colder regions with recent star formation. The velocity dispersion of the bulge (100\\,\\kms) implies in a black hole mass of $M_{BH}=8.3^{+3.2}_{-2.2}\\times10^6$\\,M$_\\odot$.  Emission-line flux distributions are most extended along PA$=27/153^\\circ$, reaching at least 450\\,pc from the nucleus and following the orientation observed in previous optical emission-line \\oiii\\ imaging and radio jet. The molecular \\h2\\ gas has an excitation temperature $T_{\\rm exc}\\approx2300$\\,K and its emission is dominated by thermal processes, mainly due to X-ray heating by the active nucleus, with a possible small contribution from shocks produced by the radio jet. The \\feii\\ excitation has a larger contribution from shocks produced by the radio jet, as evidenced by the line-ratio maps and velocity dispersion map, which show spatial correlation with the radio structures. The coronal lines are resolved, extending up to $\\approx$\\,150\\,pc and are also slightly more extended along PA$=27/153^\\circ$.  The gaseous kinematics shows two components, one due to gas located in the galaxy plane, in similar rotation to that of the stars and another in outflow, which is oriented close to the plane of the sky, thus extending to high latitudes, as the galaxy plane is inclined by $\\approx$\\,45$^\\circ$ relative to the plane of the sky. The gas rotating in the plane dominates the \\h2\\ and \\pb\\ emission, while the gas in outflow is observed predominantly in  \\feii\\ emission. The \\feii\\  emission is originated in gas being pushed by the radio jet, which destroys dust grains releasing the Fe. From the outflow velocities and implied geometry, we estimate an outflow mass rate of  $\\dot{M}_{\\rm out}\\approx6~{\\rm M_\\odot\\, yr^{-1}}$ for the ionised gas and a kinetic power for the outflow of $\\dot{E} \\approx 2.3\\times10^{41}$ erg\\,s$^{-1}\\,\\approx 0.15\\times L_{\\rm bol}$.  The distinct flux distributions and kinematics of the \\h2\\ and \\feii\\ emitting gas, with the former more restricted to the plane of the galaxy, and the later tracing the outflows related to radio jets is a common characteristic of the 6 Seyfert galaxies (ESO\\,428-G14, NGC\\,4051, NGC\\,7582, NGC\\,4151, Mrk\\,1066 and now Mrk\\,1157) we have studied so far using similar 2D observations, and other 2 (Circinus and NGC\\,2110) using long-slit observations. We conclude that the \\h2\\ emission surrounding the nucleus in the galaxy plane is a tracer of the gas feeding to the active nucleus while the \\feii\\ emission is a tracer of its feedback. ",
        "introduction": "Recent resolved imaging and spectroscopic studies of the central regions of  nearby active galaxies have been allowing us to probe the feeding and feedback mechanisms of their central engines. The study of the of the ionized gas of the Narrow-Line Region (NLR) of active galaxies allows to investigate how the radiation and mass outflows from the nucleus interact with the circumnuclear gas, affecting its kinematics and excitation \\citep[e.g.][]{fischer11,fischer10,crenshaw10a,crenshaw10b,crenshaw09,crenshaw07,kraemer09,holt06,veilleux05,schmitt96,veilleux97,wilson93}. On the other hand, the study of molecular and low ionization gas can tell us about the feeding of the Active Galactic Nucleus (AGN) of these galaxies by the mapping of inflows and estimating mass inflow rates towards the centre \\citep[e.g.][]{vandeVen10,sanchez09,sb07,fathi06,mundell99}. However, most of the above studies are based on optical observations, which are affected by dust obscuration \\citep{ferruit00,mulchaey96,mulchaey96b}, a problem that can be alleviated by using infrared lines to map the NLR emission.   Since 2006, we have been observing the inner few hundred parsecs of nearby active galaxies with Integral Field Spectrographs at the Gemini Observatory, with the goal of mapping both inflows and outflows around nearby active galactic nuclei and try to constrain the mass flow rates. In the near-infrared (hereafter near-IR) our main findings have been that the molecular (\\h2) and ionised gases present distinct flux distributions and kinematics. The \\h2\\ emission gas is usually restricted to the plane of the galaxy, while the ionised gas extends also to high latitudes and is associated with the radio emission \\citep{riffel06,n4051,n7582,mrk1066a,mrk1066c,sb09,sb10}. The \\h2\\ kinematics is usually dominated by rotation in the plane, including in some cases streaming motions towards the nucleus, while the kinematics of the ionised gas, and in particular of the \\feii\\ emitting gas, shows, in addition, a strong outflowing component associated with radio jets from the Active Galactic Nucleus (AGN). A previous study using long-slit near-IR data \\citep{sb99} of the Seyfert galaxies Circinus and NGC\\,2110, showing higher velocity dispersion and more disturbed kinematics for the \\feii\\ emission as compared with the \\h2\\ emission also supports this scenario. These results so far suggest that the molecular gas can be considered a tracer of the feeding of the AGN and the ionised gas a tracer of its feedback. Nevertheless, our 2D observations -- which provide a complete coverage of the kinematics -- comprise so far only half a dozen galaxies, and, in order to compensate for uncertain projection and filling factors, more objects need to be studied in detail, in order to allow also the quantification of the mass inflow and outflow rates.  In this work, we present the gaseous flux distribution and kinematics of the inner 450 pc of another active galaxy, the Seyfert~2 galaxy Mrk\\,1157 (NGC\\,591), for which we were able to measure also the stellar kinematics, allowing its comparison with the gas kinematics, which can be thus better constrained. This is only the second case in which we could measure the stellar kinematics, the first being Mrk\\,1066 \\citep{mrk1066c}. Mrk\\,1157 was selected for this study because: (i) it presents strong near-IR emission lines \\citep[e.g][]{nagar99,rogerio06}, allowing the mapping of the gaseous distribution and kinematics; (ii) it has extended radio emission, allowing the investigation of the role of the radio jet; (iii) the K-band CO absorption band heads have been observed in previous near-IR spectra \\citep{rogerio06} allowing the measurement of the stellar kinematics.  Mrk\\,1157 is an early-type barred spiral galaxy (SB0/a), located at a distance $d=61.1$\\,Mpc, for which 1\\arcsec\\ corresponds to 296\\,pc at the galaxy. Radio-continuum images at 3.6~cm and 6~cm show a radio double with extension of 1\\farcs2  oriented along the position angle PA=153$^\\circ$ \\citep{nagar99,ulvestad89}. Optical images show that the \\oiii\\ emission extends up to 4\\arcsec\\ following the orientation of the radio jet, while the \\ha\\ emission extends up to 20~\\arcsec\\ along the east-west direction \\citep{mulchaey96}. The near-IR nuclear spectrum of Mrk\\,1157 presents strong emission lines of [S\\,{\\sc iii}], He\\,{\\sc i}, H\\,{\\sc i}, \\feii\\ and \\h2\\ and the presence of high ionisation species as [Si\\,{\\sc vi}], [S\\,{\\sc viii}] and [Si\\,{\\sc x}] \\citep[e.g.][]{rogerio06,veilleux97}. Optical and near-IR spectra show no evidence of broad components  in the profiles of permitted emission lines \\citep[e.g.][]{rogerio06,veilleux97}, but spectropolarimetric observations reveal the presence of this component in \\ha\\ and \\hb\\ \\citep{moran00}. The near-IR nuclear spectrum of Mrk\\,1157 also shows absorption lines, with the CO absorptions in the H- and K-bands being the most prominent ones and including the detection of the CN absorption at 1.1$\\mu$m -- a signature of the presence of intermediate age stars \\citep{rogerio06,rogerio07}.  This paper is organised as follows. In Sec.~\\ref{obs} we describe the observations and data reduction procedures. The results are presented in Sec.~\\ref{results} and discussed in Sec.~\\ref{discussion}. We present our conclusions in Sec.~\\ref{conclusions}. ",
        "conclusions": "We have analysed two-dimensional near-IR $J-$ and $K_l-$bands spectra from the inner $\\approx$\\,450\\,pc radius of the Seyfert 2 galaxy Mrk\\,1157 obtained with the Gemini NIFS instrument at a spatial resolution of $\\approx$35~pc (0\\farcs12) and velocity resolution of $\\approx$\\,35$-$45\\,\\kms. We have mapped the emission-line flux distributions and ratios, as well as the stellar and gaseous kinematics, in the molecular and ionised emitting gas. The main conclusions of this work are:  \\begin{itemize}  \\item The stellar velocity field is  well reproduced by a model in which the stars follow circular orbits in a Plummer gravitational potential, with a major axis at PA=112$^{\\circ}\\pm5^{\\circ}$. The near side of the galaxy is the north--north-east;  \\item The stellar velocity dispersion presents a partial ring (radius $\\approx$\\,250\\,pc) of low values (50$-$60\\,\\kms) surrounded by higher $\\sigma_*$ values from the bulge stars. This low-$\\sigma_*$ ring is interpreted as being originated in colder regions with more recent star formation which still keep the lower velocity dispersion of the gas from which the stars have formed;  \\item The stellar velocity dispersion of the bulge ($\\approx$\\,100\\,\\kms) implies a black hole mass of $M_{BH}=8.3^{+3.2}_{-2.2}\\times10^6$\\,M$_\\odot$;  \\item Emission-line flux distributions of molecular hydrogen and low-ionisation gas are most extended along PA$=27/153^\\circ$ in agreement with previous optical \\oiii\\ imaging and following the radio jet. This emission extends to at least $\\approx$\\,450\\,pc (1\\farcs5) from the nucleus. The coronal lines of [Ca\\,{\\sc viii}] and [Si\\,{\\sc vii}] are resolved extending up to $\\approx$\\,150\\,pc (0\\farcs5) and are slightly more extended also along PA$=27/153^\\circ$;  \\item The reddening map obtained via the \\pb/\\br~line ratio presents typical values of $E(B-V)=0.5$, with the highest values of up to $E(B-V)=1.8$ seen at the nucleus;  \\item  The \\h2\\ gas is excited by thermal processes and has an excitation temperature $T_{\\rm exc}\\approx2300$\\,K. Its emission is mainly due to X-ray excitation from the central AGN, but a small contribution from shocks produced by the radio jet cannot  be discarded;  \\item The \\feii\\ emission shows flux and velocity dispersion enhancements  due to shocks produced by the radio jet which destroys the dust grains in its way out from the AGN, releasing the Fe;  \\item From the comparison between the stellar and gaseous velocity fields, as well as from channel maps along the emission line profiles, we conclude that the gas has two kinematic components: (i) the first is due to gas located in the galaxy  plane, in similar rotation to that characterizing the stellar motion; this component dominates the \\h2\\ emission; (ii) the second component is in outflow, which is oriented along the radio jet at PA=153$^\\circ$ with the near (blueshifted) side to the north-west, but is launched close to the plane of the sky, making an angle less than $\\approx$\\,30$^\\circ$ with the galaxy plane; this component extends to high latitudes and dominates the \\feii\\ emission;  \\item The ionised gas mass outflow rate through a cross section with radius $\\approx$\\,100\\,pc located at a distance of $\\approx$180\\,pc from the nucleus is $\\dot{M}_{\\rm out}\\approx6~{\\rm M_\\odot\\, yr^{-1}}$;  \\item Revised mass outflow rates from our previous studies of 8 other Seyfert galaxies, using the better constrained filling factors and gas densities we have determined in the present study, are in the range 0.13--6.25 M$\\odot$yr$^{-1}$, thus similar to the case of Mrk\\,1157. These rates are much larger than the accretion rates to the AGN, implying that most of the NLR outflow mass originates in the interstellar medium of the galaxy, which is pushed by a less massive nuclear outflow.  \\item The coronal gas, although being less extended, has similar kinematics to that of the low-ionisation gas, suggesting an origin in the inner NLR.  \\end{itemize}  Our final statement, which applies to all Seyfert galaxies we have recently  observed with NIFS, is that the \\h2\\ and the \\feii\\ emitting gas sample different flux distributions and kinematics, indicating a distinct origin. The molecular gas is in rotation in the plane of the galaxy, where is the source of fuel for  the super-massive black hole. The \\feii\\ emission, on the other hand, has higher velocity dispersions and extends to high latitudes, mapping mostly the AGN outflows."
    },
    "1107/1107.6037_arXiv.txt": {
        "abstract": "Collisionless plasma instabilities are fundamental in magnetic field generation in astrophysical scenarios, but their role has been addressed in scenarios where velocity shear is absent. In this work we show that velocity shears must be considered when studying realistic astrophysical scenarios, since these trigger the collisionless Kelvin-Helmholtz instability (KHI). We present the first self-consistent three-dimensional (3D) particle-in-cell (PIC) simulations of the KHI in conditions relevant for unmagnetized relativistic outflows with velocity shear, such as active galactic nuclei (AGN) and gamma-ray bursts (GRBs). We show the generation of a strong large-scale DC magnetic field, which extends over the entire shear-surface, reaching thicknesses of a few tens of electron skin depths, and persisting on time-scales much longer than the electron time scale. This DC magnetic field is not captured by MHD models since it arises from intrinsically kinetic effects. Our results indicate that the KHI can generate intense magnetic fields yielding equipartition values up to  $\\epsilon_B/\\epsilon_p \\simeq 10^{-3}-10^{-2}$ in the electron time-scale. The KHI-induced magnetic fields have a characteristic structure that will lead to a distinct radiation signature, and can seed the turbulent dynamo amplification process. The dynamics of the KHI are relevant for non-thermal radiation modeling and can also have a strong impact on the formation of relativistic shocks in presence of velocity shears. ",
        "introduction": "Magnetic field generation in astrophysical scenarios, such as AGN and GRBs, is not fully understood \\citep{colgate01}. While such phenomena are of fundamental interest, they are also closely related to open questions such non-thermal radiation emission and cosmic ray acceleration \\citep{bhattacharjee00}. Recently, collisionless plasma effects have been proposed as candidate mechanisms for magnetic field generation \\citep{gruzinovwaxman99, medvedev99}, mediating the formation of relativistic shocks via the Weibel instability \\citep{weibel59,silva03}. The KHI \\citep{dangelo65,gruzinov08,macfadyen09} should also be considered since it is capable of generating large-scale magnetic fields in the presence of strong velocity shears, which naturally originate in energetic matter outbursts in AGN and GRBs, and which are also present whenever conditions for the formation of relativistic shocks exist. These large-scale fields may also be further amplified by the magnetic-dynamo effect \\citep{gruzinov08,macfadyen09}.  Recent kinetic simulations have focused on magnetic field generation via electromagnetic plasma instabilities in unmagnetized flows without velocity shears. 3D PIC simulations of Weibel turbulence \\citep{silva03,fonseca03,frederiksen04,nishikawa05} have demonstrated subequipartition magnetic field generation. The Weibel instability has been shown to be critical in mediating relativistic shocks \\citep{spitkovsky08,martins09}, where a Fermi-like particle acceleration process has also been identified. These works have neglected the role of velocity shear in the flow, which are an alternative mechanism to generate sub-equipartition magnetic fields in relativistic outflows \\citep{gruzinov08}. Furthermore, a shear flow upstream of a shock can lead to density inhomogeneities via the KHI which may constitute important scattering sites for particle acceleration. In \\cite{macfadyen09}, 3D magnetohydrodynamic (MHD) simulations of KH turbulence are discussed, observing magnetic field amplification due to the KHI. In fact, the KHI setup is routinely used to test MHD models \\citep{mignone09, beckwith11}, but the connection between the MHD description and the fully kinetic picture is still missing \\citep{gruzinov11}. However, the KHI contains intrinsically kinetic features, and so far 3D fully kinetic ab initio simulations have not been reported.  In this Letter, we present the first self-consistent 3D PIC simulations of the KHI for both subrelativistic and relativistic scenarios of shearing unmagnetized electron-proton plasma clouds. We show that the KHI contains important kinetic features which are not captured in previous 3D MHD simulations \\citep{keppens99,macfadyen09,mignone09,beckwith11}, namely the transverse KHI dynamics, and the generation of a large-scale DC magnetic field at the shear region. This large-scale field can reach mG levels for typical parameters of the interaction of a relativistic flow with the interstellar medium (ISM). Furthermore, our generalization of the KHI linear theory \\citep{gruzinov08} to include arbitrary density jumps between the shearing flows allows us to conclude that the onset of the instability is robust to this asymmetry. The KHI can therefore operate in shears within the ejecta (similar density flows) and also between the ejecta and the surrounding ISM (relative density ratios of $1-10$), and it will likely operate at the same level (or stronger) than the Weibel instability, even for moderate velocity shears. In Section 2, we explore the density jumps effect on the behavior and features of the instability. The 3D PIC simulation results are presented and discussed in detail in Section 3. In Section 4, we discuss the saturation levels of the magnetic field, and conclusions are drawn in Section 5. ",
        "conclusions": "In this Letter, we present the first self-consistent 3D PIC simulations of the KHI in unmagnetized electron-proton plasmas and analyze their evolution on the electron time-scale. Our results show that the multidimensional physics of the KHI is extremely rich and that kinetic effects play an important role, in particular, in the transverse dynamics of the KHI (which is dominant over the longitudinal KHI dynamics in relativistic shears), and in the generation of a strong large-scale DC magnetic field.  The transverse dynamics of the KHI consists of a Weibel-like electron bunching process, leading to the formation of electron current filaments which are then accelerated across the shear surface, forming finger-like structures. At the electron saturation time-scale, the magnetic field has evolved to a large-scale DC field structure that extends over the entire shear-surface, reaching thicknesses of a few tens  of electron skin depths, and persisting on time-scales much longer than the electron time-scale. This field structure is not captured by MHD or other fluid models, and will lead to a distinct radiation signature. We measure maximum equipartition values of $\\epsilon_B/\\epsilon_p \\simeq 2 \\times 10^{-3}$ for the subrelativistic scenario, and $\\epsilon_B/\\epsilon_p \\simeq 7 \\times 10^{-3}$ for the relativistic scenario. These equipartition values, which are typically treated as a free parameter in radiation modeling, match the values inferred from GRB afterglows by \\citep{panaitescu02}. Moreover, the onset of the KHI is robust to density asymmetries making it ubiquitous in astrophysical settings. The KHI may operate in $n_+/n_- \\approx 1$ regimes, relevant in GRB internal shocks, and also in $n_+/n_- \\gg 1$ regimes, which are important in external shocks. Future work will address the impact of the KHI in the formation of relativistic shocks in the presence of velocity shears.   \\vspace{0.2cm}  \\small  This work was partially supported by the European Research Council ($\\mathrm{ERC-2010-AdG}$ Grant 267841) and FCT (Portugal) grants SFRH/BD/75558/2010, SFRH/BPD/75462/2010, and PTDC/FIS/111720/2009. We would like to acknowledge the assistance of high performance computing resources (Tier-0) provided by PRACE on Jugene based in Germany. Simulations were performed at the IST cluster (Lisbon, Portugal), and the Jugene supercomputer (Germany)."
    },
    "1107/1107.1798_arXiv.txt": {
        "abstract": "The first observations of seismic responses to solar flares were carried out using time-distance (TD) and holography techniques applied to SOHO/MDI Dopplergrams obtained from space and un-affected by terrestrial atmospheric disturbances. However, the ground-based network GONG is potentially a very valuable source of sunquake observations, especially in cases where space observations are unavailable. In this paper we present updated technique for pre-processing of  GONG observations for application of subjacent vantage holography. Using this method and TD diagrams we investigate several sunquakes observed in association with M and X-class solar flares and compare the outcomes with those reported earlier using MDI data. In both GONG and MDI datasets, for the first time, we also detect the TD ridge associated with the September 9, 2001 flare. Our results show reassuringly positive identification of sunquakes from GONG data that can provide further information about the physics of seismic processes associated with solar flares. ",
        "introduction": "Discovery of sunquakes associated with solar flares \\citep{kz1998,Donea1999} has opened a new era in the  investigation of energy and momentum transport mechanisms from the upper atmosphere to the photosphere and beneath, uncovering structure of these spectacular events. Sunquakes, seen as circular or elliptical waves - ripples, propagating outward from impulsive hard X-ray (HXR) solar flares along the solar surface, appear 20-60 minutes after the flare onsets. The surface ripples are also associated with strong downward shocks preceding these ripples with close (1-4 minutes) temporal correlation with the start of HXR flares, indicating that high energy particles play some role in initiation of sunquakes \\citep{Zharkova07, Martinez2008a}.  Even though every flare is expected to inject particle beams of one kind or another into a flaring atmosphere, inducing either shocks or magnetic impulses, not many of them have { recorded measurable} signatures of seismic activity. Initially only X-class flares were considered as candidates for producing sunquakes. The first flare detection used time-distance diagrams and reported well distinguished ripples emanating from the center of the location of a hard X-ray source in the X1.1 flare 9 July 1996 \\citep{kz1998, Donea1999}. It was followed by detection of quakes associated with two extremely powerful solar flares of class X17 and X10 which erupted in NOAA Active region 10486 on October 28 and 29, 2003. These two flares, known as the Halloween 2003 flares, were extensively investigated in \\cite{DL2005} by applying subjacent vantage acoustic holography to SOHO/MDI Dopplergram data. The  acoustic signatures were also shown to co-align with the hard X-ray signatures and GONG intensity observations revealed significant radiative emission with a sudden onset in the compact region encompassing the acoustic signature.  Further analysis of the 29 October 2003 quake was presented in \\citet{LD2008}, where the authors proposed a new method for correcting { intensity data recorded by the Global Oscillation Network Group (GONG)}, that allowed comparison of acoustic kernels with white-light traces of the flare. Later \\citet{VZharkova2006, K2006, Zharkova07} investigated the Halloween flare of 28 October 2003 using the time-distance diagram method applied to SOHO/MDI data and detected three distinct seismic sources that coincided with the holographic sources from \\cite{DL2005}. { For the 29 October 2003 flare there were no time-distance ridges found by the authors of \\citet{Zharkova07} when they analyzed MDI dopplergrams, however one was later reported in \\cite{K2006}}. The first M-class flare in which seismic signatures were detected  by means of acoustic holography using Doppler velocity data from SOHO/MDI instrument was the flare that occurred in NOAA Active Region 9608 on September 9, 2001 \\citep{Donea2006}.  Later the list of acoustically active flares was significantly extended. \\citet{BI05} reported another six such flares, adding another four in \\citet{DoneaEtAl2006}. In fact, during the period of observations with SOHO/MDI instrument there have been 17 flares showing signs of acoustic activity possibly related to sunquakes\\footnote{see \\url{http://users.monash.edu.au/$\\sim$dionescu/sunquakes/sunquakes.html}}. This expanding list motivated researchers to look further and to explore the data from the  ground-based GONG observatories, which offers better coverage of helioseismic data compared to  SOHO/MDI.  In the unusually quiet solar minimum between cycles 23 and 24 more attention was paid to each new flare occurring on the Sun, one of which was the flare of 14 December 2006. { The latter was observed only by the GONG instruments and revealed some noticeable seismic signatures} in both time-distance ridges and egression powers \\citep{Matthews10}.   In order to validate { these findings, a comparison is required of the signatures of sunquakes derived for both the time-distance and holographic techniques} from GONG data with those from SOHO/MDI for a number flares with distinct seismic signatures. Such a comparison will allow us  { to understand the differences in appearances and} to derive recommendations for the reliable detection of sunquakes from GONG data.  In this study we consider three acoustically active flares that have the luxury of  helioseismic observations available from both the GONG and the SOHO/MDI instrument. { The available Dopplergrams (GONG and MDI) are used to  analyze the acoustic signatures of the flares by using both the time-distance diagram technique \\citep[TD method;][]{kz1998} and acoustic holography \\citep[e.g.][]{BL1999, LB2000}.} The description of data and additional corrections for the technique applied to GONG data are presented in section \\ref{descrip}, the results of the comparison are described in section \\ref{valid} and conclusions with recommendations are drawn in section \\ref{conc}. ",
        "conclusions": "\\label{conc} In this study we have compared egression power maps and time-distance diagrams derived from GONG and MDI data. SOHO MDI and GONG velocity datasets were used for three flares: M-class September 9, 2001 (Figures \\ref{fig:20010909_egression}-\\ref{fig:20010909_holo}, \\ref{fig:20010909_TD}), X-class October 28, 2003 (Figures \\ref{fig:20031028_holo} and \\ref{fig:20031028_TD}), X-class October 29, 2003 (Figure \\ref{fig:20031029_holo}).  Reassuringly, the egression power map snapshots show seismic signatures common to both instruments for all flares. These signatures display an excellent agreement between the two instruments in terms of their time and location. We note, however, that even after the pre-processing, as outlined in Section \\ref{sec:addGONG}, GONG egression measurements remain relatively noisy. This leads to important differences from the MDI produced egression maps, which are only partially compensated by increasing the pupil sizes when working with GONG data. One such difference is the apparent variance in shape and strength of the detected acoustic kernels as seen by these two instruments. Another is the relative weakness of the suppression of acoustic sources below the sunspot. Since solar quakes are often located in the sunspot, this means that identifying such seismic signatures in GONG egression measurements is a harder task due to its lower signal to noise ratio over the sunspot region.  One such method of verification is the computation of the time-distance diagram. As we have demonstrated for September 9, 2001 and October 28, 2003 flares, such diagrams computed from GONG velocity data can present the additional evidence of the quake. Figures \\ref{fig:20010909_TD} and \\ref{fig:20031028_TD} clearly show that, in spite of being less sensitive, the GONG data can respond to the time-distance analysis producing noticeable ridges similar to those observed from the higher-quality MDI data. Results of the comparison with MDI time-distance measurements have again revealed an excellent spatial agreement between the two instruments in terms of the time-distance source location. Additionally, the fact that the locations of the sources observed with the GONG time-distance diagrams coincide with the acoustic kernels deduced from the GONG egression snapshots confirms that with these different techniques one observes the same events - seismic signatures induced by solar flares. Therefore, we conclude that the GONG data can respond to time-distance analysis. Obviously, due to the characteristic noise, not every quake can be expected to produce the time-distance ridge in GONG diagrams, but  if a ridge is seen in the GONG ground-based data, one can expect that it will also be observed by using the higher-quality satellite MDI or HMI data.  We believe, the results of this study show that quake detection based on helioseismic reduction of GONG observations is possible. However, as the data are subject to atmospheric smearing and other related instrumental effects, GONG observations clearly have less intrinsic sensitivity than the space-borne observations. A useful prospective object for further study might be the quantitative comparison of intrinsic sensitivities of ground- and space-based helioseismic observations under various seeing conditions. Nonetheless these results should allow us to add to the list of known sun-quakes by investigating the known flares in the Solar Cycle 23 using GONG data when MDI observations were not available. This will provide further information about the physics of seismic processes associated with solar flares."
    },
    "1107/1107.3562_arXiv.txt": {
        "abstract": "We consider a scenario where supermassive black holes form through direct accumulation of gas at the centre of proto-galaxies. In the first stage, the accumulated gas forms a super-massive star whose core collapses when the nuclear fuel is exhausted, forming a black hole of $M_{\\rm BH} \\approx 100 M_{\\sun}$. As the black hole starts accreting, it inflates the surrounding dense gas into an almost hydrostatic self-gravitating envelope, with at least $10-100$ times the mass of the hole. We find that these ``quasistars\" suffer extremely high rates of mass loss through winds from their envelopes, in analogy to very massive stars such as $\\eta$-Carinae. Only for envelope masses greater than $2.8 \\times 10^{5} (M_{\\rm BH}/100 M_{\\sun})^{9/11}$ is the envelope evaporation time-scale longer than the accretion time-scale of the black hole. This relation thus constitutes a ``threshold growth line\" above which quasistars can grow their internal black holes. Accretion rates can be $10$ to $100$ times the Eddington rate. The quasistars born in this ``growth region\" with $10^7-10^8 M_{\\sun}$  can grow black holes with masses between $10^4$ to $10^5 M_{\\sun}$, before crossing the threshold growth line and dispersing their envelopes in less than $10^4$ yr. This scenario therefore predicts that massive black hole seeds can be found only in dark matter halos with total masses larger than about $10^9 M_{\\sun}$, which can provide sufficiently high accretion rates to form such massive quasistars. ",
        "introduction": "\\label{sec:intro} The formation of supermassive black holes (SMBHs) requires the inflow of $ > 10^6 M_{\\sun}$ of gas from galactic scales into a region much smaller than a parsec. The presence of bright ($\\sim 10^{47}$ erg s$^{-1}$) quasars at $z >6$ \\citep{Fan2001} suggests that, at least in those cases,  this inflow must have started at an earlier epoch and proceeded at sufficiently high rate to have allowed the assemblage of black holes (BHs) of $10^9 M_{\\sun}$, in less than a Gyr.  In proto-galactic dark matter halos with mass $> 10^9 M_{\\sun}$, gas can be funnelled towards the centre at a rate of $> M_{\\sun}$ yr$^{-1}$.  If substantial fragmentation of gas into stars can be avoided \\citep[e.g.,][]{Wise2008}, the collection of gas at such high rates can lead to the formation of a light ($\\lesssim 100 M_{\\sun}$) black hole (BH) via direct collapse of gas or through the intermediate stage a supermassive star, whose core collapses to form a BH \\citep{Begelman2010}. The feedback from this ``embryo\" BH becomes important and it modifies the structure of the pregalactic discs in its innermost $\\approx 100$ AU. The BH luminosity inflates the highly opaque gas in its vicinity into a pressure supported envelope, at least $10-100$ times more massive than the BH itself. While the envelope is being fed by the proto-galactic disc, the BH accretes from the envelope. This structure is dubbed a {\\it quasistar}, since it resembles a (scaled-up) red giant in structure, though it is powered by accretion into a central BH \\citep[][hereafter BRA08]{Begelman2006,Begelman2008}.   BRA08 found that for each black hole mass $M_{\\rm BH}$, there is a minimum envelope mass $M_*$ below which no hydrostatic solutions for the whole envelope can be obtained. This line corresponds to modest $M_*/M_{\\rm BH}$ ratios of a few tens. Above this line, the BH accretes at the Eddington rate for the whole mass ($M_{\\rm T} = M_{\\rm BH}+M_*$). The BH can accrete at a super-Eddington rate (for the BH mass), because the energy released by the accretion is transported outwards through convection (and not by radiative transfer).  After a few Myr, the BH reaches a maximum mass of $\\sim 10^4 M_{\\sun}$,  which provides a luminosity that unbinds the surrounding envelope. This physical picture has been confirmed by numerical simulations (Ball et al. 2011, hereafter BTZE11). Accretion will then proceed at a more ``modest\" rate through the proto-galactic disc. Accretion and mergers with other holes allow these ``BH seeds\" to evolve over cosmic times into SMBHs.  An important characteristic of quasistars is that they are loosely bound, since they are radiation pressure dominated. A luminosity that is close to the Eddington limit therefore has the potential to drive strong radiation-driven outflows.  The presence of an outflow, in turn, would modify the quasistar's internal structure, the BH accretion rate and the final mass of the BH seed that can be produced. Quantitatively, however, the extent and consequences of such winds are currently not known.  In this paper, we aim to address the above questions. We calculate approximate solutions for a hydrostatic massive envelope powered by accretion, allowing for mass loss from its surface. We find that the presence of a wind allows for the existence of hydrostatic solutions for the envelope, even if the system is super-Eddington (for the total mass). With the wind, there is a competition between accretion and evaporation, and the BH can be effectively fed and increase its mass only when the envelope evaporation is sufficiently weak. This happens for very massive envelopes $> 10^6 M_{\\sun}$, whose wind reaches a maximum mass loss, driven by all the available luminosity.  We find that under the above conditions, a BH seed of $> 10^3$ solar masses can be assembled in less than $10^4$ yr. In particular, to obtain BH seeds of $10^4-10^5 M_{\\sun}$, an envelope of mass $>10^7 M_{\\sun}$ is required. We will show that, as a consequence, the present scenario predicts the presence of massive BH seeds at high redshift ($z \\gtrsim 10$) {\\em only} in halos of least $\\approx {\\rm a \\ few} \\ 10^{9} \\,M_{\\sun}$.   The paper is organized as follows. In \\S\\ref{sec:analytics}, we describe our model and carry out analytical estimates. In \\S\\ref{sec:numerical}, we describe the numerical solution for the structure of super-Eddington quasistars with winds, and we present our results in \\S\\ref{sec:results}. We discuss the implications of our findings to the modeling of supermassive BH formation in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In the present work, we considered the state and evolution of quasistars, where a massive envelope enshrouds a black hole. We found that an important characteristic of these objects is the continuum driven winds which they accelerate. This gives rise to an evaporation strip in the envelope mass -- BH mass parameter space, where quasistars evaporate. BH growth takes place only for large enough $M_*$ for which the wind mass loss is limited by photon tired winds.  We then considered two evolutionary scenarios. In the first scenario, the quasistar is initially formed above the evaporation strip, by first having a large gas cloud collapse, then form an internal black hole, and subsequently evolve. In the second scenario, a seed BH is first formed, and then it accretes gas at large rates. We found that only the first scenario can generate large seed BHs. This is because once a BH is formed, unrealistically high mass accretion rates are required to evolve the quasistar out of the evaporation strip.  These results put strong constraints on the dark matter halos in which massive SMBH seeds can form. Our first scenario, on which we now concentrate, follows the picture put forth by \\cite{Begelman2010}, whereby the progenitors of Quasistars are super massive stars that formed as the consequence of the high infall rate of hydrogen-cooled gas at the centre of dark matter halos. The life-time of these stars is set by the thermonuclear timescale for burning their hydrogen core.  At the Eddington limit, it  is $\\sim 2$ Myr, independent of mass. After that, the core may collapse into a BH. Once the BH starts accreting from the envelope, the feedback from the released luminosity sets up the structure we have investigated in this paper.  Our results constrain the initial mass of the quasistar, and thus of the supermassive star, that can grow massive SMBH seeds. A quasistar with an initial mass $\\approx 10^6 M_{\\sun}$ (or $\\approx 10^7 M_{\\sun}$) can form a BH of $\\approx 10^3 M_{\\sun}$ (or $ \\approx 10^4 M_{\\sun}$ respectively, as can be see in Fig.\\ 4). This initial mass should be accumulated  in less than $\\sim 2$ Myr, which requires accretion rates greater than a few $M_{\\sun}$ yr$^{-1}$.   Dark matter halos accrete matter through their virial radius at a rate of \\be {\\dot M}_{\\rm DM} \\approx  4.6  \\ (z+1)^{2.5}_{10}  \\left(\\frac{M_{\\rm h}}{10^9 M_{\\sun}}\\right)^{1.14}  {\\rm M_{\\sun} \\, yr^{-1}}, \\ee  \\no \\citep{Neistein+06}, where $M_{\\rm h}$ is the mass of the dark matter halo that hosts the quasistar and $(z+1)_{10} = (z+1)/10$.  Since the timescale for dark matter and gas accretion is the same \\citep{Dekel+09}, \\be {\\dot M}_{\\rm acc} = f_{\\rm b} \\times {\\dot M}_{\\rm DM}, \\label{eq:macc_cosmo} \\ee it is also an estimate for the accretion rate that feeds the quasistar, under the assumption that almost all gas can be funnelled towards the halo centre. In eq.~\\ref{eq:macc_cosmo}, $f_{\\rm b} \\approx 0.17$ is the cosmological baryon fraction \\citep{komatsu+2010}.  In fact, it is not clear if such high accretion rates do not undergo fragmentation at parsec scales  and if so, how much gas turns into stars (e.g., \\citealt{Shlosman1989}; \\citealt{Levin2007}). This is known as the ``active galactic nuclei (AGN) fueling problem\", since we observe them shining with a luminosity that implies $\\sim 1 M_{\\sun}$ yr$^{-1}$. However if, like it must happen in AGNs, this problem is not severe \\citep{Wise2008,Begelman2009} such flows can reach the proto-galactic centre.  Under this assumption, we can thus connect the accretion rate needed in order to form a quasistar of a certain mass and at a certain redshift, with the mass of the host halo through eq.~\\ref{eq:macc_cosmo}. However, we should also assume an efficiency factor, since not all gas $f_{\\rm b} \\times M_{\\rm h}$ in the halo can be used to form a quasistar. Therefore, the minimum halo mass that can host a given quasistar with mass $M_*$ is  \\be \\frac{M_{\\rm h}}{M_{\\sun}} \\approx \\max \\left[\\frac{7 \\times 10^{8} }{(1+z)_{10}^{2.2}}\\left(M_* \\over10^6 M_{\\sun}\\right)^{0.9},6 \\times 10^8 \\frac{M_*}{10^6 M_{\\sun}}\\right], \\label{mh_min} \\ee  \\no where we assumed in the second term that no more than $1\\%$ of the total amount of gas is used for the formation of the quasistar. The first term comes directly from eq.~\\ref{eq:macc_cosmo}.  Eq.~\\ref{mh_min} implies that at $z=10$, quasistars with $M_*>10^6 M_{\\sun}$ can be found only in dark matter halos with $M_{\\rm h} \\gtrsim 10^9 M_{\\sun}$. In particular, for $z \\gtrsim 10$ ``massive\" SMBH seeds of $M_{\\rm BH} >10^4 M_{\\sun}$ need host halos with $M_{\\rm h} \\gtrsim 6 \\times 10^{9} M_{\\sun}$.  The very bright quasars observed at $z \\sim 6$ have masses of typically $10^9 M_{\\sun}$, estimated using the Eddington argument. Moreover, seeds of $10^4-10^5 M_{\\sun}$ can grow by accretion at the Eddington rate in $\\sim 0.5$ Gyr. Therefore, such SMBH seeds may have formed as late as $z\\sim 10$. At this redshift, the comoving number density of halos with $M_{\\rm h} \\approx 10^{12} M_{\\sun}$  (calulated with the \\citealt{ST99} formalism) matches the observed comoving number density of those bright quasars, $10^{-9} {\\rm Mpc}^{-3}$ \\citep{Fan2001}. Those halos are indeed above the threshold set by eq.\\ \\ref{mh_min}.  Of course, halos formed at different redshifts can lead to the observed high redshift quasars and to the observed local galaxies with SMBHs. To properly derive the consequences of eq.\\ \\ref{mh_min}, one should thus follow a ``merging tree\" evolution of halos and also of the hosted black holes. This is beyond the scope of this paper, but it will be addresses in a forthcoming work.   Finally, our results show that, unfortunately, Quasistars will not be easy to observe. They are relative rare  objects which shine at super-Eddington rates only for a short period, prior to their final evaporation."
    },
    "1107/1107.0757_arXiv.txt": {
        "abstract": "We present a large sample of over 200 integrated-light spectra of confirmed globular clusters (GCs) associated with the Sombrero (M104) galaxy taken with the DEIMOS instrument on the Keck telescope. A significant fraction of the spectra have signal-to-noise levels high enough to allow measurements of GC metallicities using the method of Brodie \\& Huchra (1990). We find a distribution of spectroscopic metallicities ranging from --2.2 $<$ [Fe/H] $<$ +0.1 that is bimodal, with peaks at [Fe/H] $\\sim$ --1.4 and --0.6. Thus the GC system of the Sombrero galaxy, like a few other galaxies now studied in detail, reveals a bimodal {\\it spectroscopic} metallicity distribution supporting the long-held belief that colour bimodality reflects two metallicity subpopulations. This further suggests that the transformation from optical colour to metallicity for old stellar populations, such as GCs, is not strongly non-linear. We also explore the radial and magnitude distribution with metallicity for GC subpopulations but small number statistics prevent any clear trends in these distributions. ",
        "introduction": "Understanding how galaxies were formed and have evolved remains one of the basic problems to be solved in astrophysics. In this context, extragalactic globular clusters (GCs) play an important role because their formation is associated with the physical processes occurring before galaxies were entirely assembled (see Brodie \\& Strader 2006 for a review).  It has been recognised for some time that massive galaxies have bimodal GC colour distributions (Ashman \\& Zepf 1993). This suggests two modes, or phases, of GC formation. Indeed, several models have been proposed to explain the different modes in terms of variations in the epoch of GC formation, stellar population properties and the timescale of galaxy assembly (Ashman \\& Zepf 1992; Zepf \\& Ashman 1993; Forbes et al. 1997; Cote et al. 1998; Beasley et al. 2002). As most GCs appear to be very old (see review by Brodie \\& Strader 2006), the bimodality in colour is normally assumed to translate directly into a bimodality in metallicity. However, doubt has recently been cast on this interpretation by some (e.g. Yoon et al. (2006; Blakeslee, Cantiello \\& Peng 2010). Depending on the degree to which the transformation of colour into metallicity is non-linear, it is possible for an instrinsically unimodal metallicity distribution to appear bimodal in colour space. If correct, this would have direct implications for the observed blue tilt (a trend for redder colours with higher luminosities in the blue GC subpopulation) and the correlation of the mean GC colour with galaxy luminosity. It would also radically change current ideas for two modes of GC formation.  As noted by Blakeslee et al. (2010), {\\it \"Very little is actually known about the detailed metallicity distributions of GCs in giant ellipticals.\"}.  Two giant ellipticals with optical spectroscopic metallicities derived for their GC systems include M49 with 47 measurements (Strader et al. 2007) and NGC 5128 with over 200 GC spectroscopic metallicities (Beasley et al. 2008). In both cases the distribution of GC spectroscopic metallicities are bimodal. The GC systems of these galaxies also have bimodal optical colour distributions, as do most but not all galaxies (see e.g. Foster et al. 2010). Interestingly, the infrared study of GCs in NGC 5128 by Spitler et al. (2008), which is more sensitive to metallicity than optical colours, also revealed colour bimodality in the 146 GCs they studied. Along similar lines, Kundu \\& Zepf (2007) used optical-infrared colours of 80 GCs in M87 to show their distribution was bimodal. Thus these giant ellipticals with optical spectra and/or infrared colours of GCs show metallicity bimodality.  Recently attempts have been made to derive spectroscopic metallicities using the infrared Calcium Triplet (CaT) lines (at 8498, 8542, and 8662 $\\rm \\AA$) for the GC systems of two giant ellipticals (Foster et al. 2010, 2011). In the case of NGC 1407, the CaT metallicity distribution of 144 GCs is better described as unimodal, whereas Cenarro et al. (2007) suggested that the distribution was bimodal based on metallicities from optical wavelength spectra of just 20 GCs. For NGC 4494, the CaT metallicity distribution again appears unimodal. Unfortunately,  there is no published work based on optical wavelength metallicities for the GC system in this galaxy. Clearly a large sample study of GC metallicities derived from both blue and infrared absorption lines is needed to resolve these issues (see Foster et al. 2010 for a complete discussion of potential factors affecting the CaT derived metallicities). In summary, it appears that at least {\\it some} elliptical galaxies reveal the presence of two metallicity subpopulations in their GC systems but the number of systems studied remains small.  For late-type spiral galaxies, only two are well-studied (i.e. the Milky Way and M31) and both reveal bimodal spectroscopic metallicity distributions for their GC systems (e.g. Zinn 1985; Barmby et al. 2000 but see also Caldwell et al. 2011). No early-type spirals have published spectra for large numbers of GCs.   The nearest spiral galaxy with a large GC system is the Sombrero galaxy (M104, NGC 4594), which lies a distance of 9.0 $\\pm$ 0.1 Mpc (see Table 4 of Spitler et al. 2006 for a summary). Although classified as an edge-on Sa galaxy, it has a bulge-to-total ratio of 0.8 (Kent 1988) and hence might  be  better described as a massive bulge plus an extended disk. This galaxy hosts some 1900 GCs that extend to a projected radius of 50 kpc (Rhode \\& Zepf 2004; Spitler et al. 2006) and one Ultra Compact Dwarf (Hau et al. 2009). A number of previous studies have obtained spectra of Sombrero GCs (Bridges et al. 1997; Larsen et al. 2002; Held et al. 2003; Bridges et al. 2007), however the number of spectra of sufficient signal-to-noise (S/N) to measure individual GC metallicities was very limited. The Keck spectra of Larsen et al. (2002) confirmed that the dozen luminous GCs they studied were all old.  In this paper we present Keck integrated spectra for over 200 globular clusters in M104. This is the largest sample of GCs in the Sombrero galaxy homogeneously analysed to date. Our main goal is to obtain radial velocities to confirm association with the Sombrero galaxy and to obtain spectroscopic metallicities. A future paper will investigate the kinematic properties of the GC system. ",
        "conclusions": "Globular clusters are widely recognised as extremely useful tracers of the formation and evolution of their host galaxies. Here we present Keck spectra for over 200 GCs with radial velocities that confirm their association with the Sombrero galaxy. For many of the GCs the spectra are of sufficient quality to derive Brodie \\& Huchra (1990) style metallicities from several indices.  We find that the GCs span a broad metallicity range from [Fe/H] = --2.2 to +0.1. This is comparable to previous estimates based on photometric metallicities. For a restricted high quality subsample of 112 GCs, we find a good correlation between individual GC metallicity and B--R colour. The resulting spectroscopic metallicity distribution is clearly bimodal with peaks at [Fe/H] $\\sim$ --1.4 and --0.6, with a statistical probability of $>$ 90 per cent. Thus despite claims that colour bimodality may not reflect metallicity bimodality in extragalactic GC systems,  it would appear that the GC system of the Sombrero galaxy joins that of some other giant elliptical galaxies in revealing spectroscopic bimodality. Furthermore we find a transformation between optical colour and metallicity that does not require a non-linear relation.  We also investigate the GC radial metallicity profile out to 30 kpc and a spectroscopic version of the colour-magnitude diagram, but our small number statistics prevent any strong statements about trends in the GC subpopulations."
    },
    "1107/1107.2752_arXiv.txt": {
        "abstract": "{% As was demonstrated in earlier studies, turbulence can result in a negative contribution to the effective mean magnetic pressure, which, in turn, can cause a large-scale instability. In this study, hydromagnetic mean-field modelling is performed for an isothermally stratified layer in the presence of a horizontal magnetic field. The negative effective magnetic pressure instability (NEMPI) is comprehensively investigated. It is shown that, if the effect of turbulence on the mean magnetic tension force vanishes, which is consistent with results from direct numerical simulations of forced turbulence, the fastest growing eigenmodes of NEMPI are two-dimensional. The growth rate is found to depend on a parameter $\\betastar$ characterizing the turbulent contribution of the effective mean magnetic pressure for moderately strong mean magnetic fields. A fit formula is proposed that gives the growth rate as a function of turbulent kinematic viscosity, turbulent magnetic diffusivity, the density scale height, and the parameter $\\betastar$. The strength of the imposed magnetic field does not explicitly enter provided the location of the vertical boundaries are chosen such that the maximum of the eigenmode of NEMPI fits into the domain. The formation of sunspots and solar active regions is discussed as possible applications of NEMPI. }  \\begin{document} ",
        "introduction": "The concept of turbulent viscosity is often used in astrophysical and other applications in recognition of the fact that the microscopic viscosity is far too small to be relevant on the length scales under consideration. Turbulent viscosity is the simplest parameterization of the Reynolds stress tensor, $\\overline{u_i u_j}$, where $\\uu=\\UU-\\meanUU$ is the velocity fluctuation about a suitably defined average, denoted here by an overbar. Turbulent viscosity is by far not the only contribution to the Reynolds stress tensor. In addition to hydrodynamic contributions such as the $\\Lambda$ effect (R\\\"udiger 1980, 1989), which is relevant to explaining stellar differential rotation (R\\\"udiger \\& Hollerbach 2004), and the anisotropic kinetic alpha effect (Frisch et al.\\ 1987), which provides an important test case in mean-field hydrodynamics (Brandenburg \\& von Rekowski 2001; Courvoisier et al.\\ 2010), there are magnetic contributions as well. One can think of them as a magnetic feedback on the hydrodynamic stress tensor (R\\\"adler 1974; R\\\"udiger 1974) or, especially when magnetic fluctuations are also considered, as a mean-field contribution to the turbulent Lorentz force.  Work by Roberts \\& Soward (1975) and R\\\"{u}diger et al.\\ (1986, 2012) using a quasi-linear approach suggests that the total magnetic tension force (which includes the effects of fluctuations) is reduced in the presence of a mean magnetic field and might formally even change sign for larger magnetic Reynolds numbers, but this would then be beyond the validity of their approximation. Using the spectral $\\tau$ approach, Kleeorin et al.\\ (1989, 1990) do indeed find this reversal of the sign of the total magnetic tension force. In addition, they find a reversal of the sign of the effective magnetic {\\it pressure} term; see also Kleeorin \\& Rogachevskii (1994) and Kleeorin et al.\\ (1993, 1996). Rogachevskii \\& Kleeorin (2007) argue that, in a stratified medium, this can lead to the formation of large-scale magnetic flux structures and perhaps even sunspots -- or at least active regions.  Recently, direct numerical simulations (DNS) of both unstratified and stratified forced turbulence (Brandenburg et al.\\ 2010, 2012; hereafter referred to as BKR and BKKR, respectively) have substantiated this idea and have demonstrated that the effective magnetic pressure can indeed change sign. Similar results have now also been obtained for turbulent convection (K\\\"apyl\\\"a et al.\\ 2012). In addition, these papers give results of mean-field calculations illustrating that there is a negative effective magnetic pressure instability (hereafter referred to as NEMPI) when there is sufficient density stratification.  NEMPI is a convective type instability related to the interchange instability in plasmas (Tserkovnikov 1960; Newcomb 1961; Priest 1982) and the magnetic buoyancy instability in the astrophysical context (Parker 1966). The free energy in interchange and magnetic buoyancy instabilities is drawn from the gravitational field, while in NEMPI it is provided by the small-scale turbulence.  The mechanism of NEMPI works even under isothermal conditions when entropy evolution is ignored and an isothermal equation of state is used. This has been shown using corresponding mean-field calculations (BKKR). With this reduction to the most elementary aspects of the instability, it has recently been possible to verify the existence of NEMPI also in DNS (Brandenburg et al.\\ 2011, hereafter referred to as BKKMR). This has been a major step forward, because now there is no doubt that one is pursuing a real effect and not just one that works only in the world of mean-field models. Essential to the paper of BKKMR has been a finding from an earlier version of the present one that only two-dimensional mean-field structures are excited. This property allowed meaningful averaging along the direction of the imposed field, making the identification of flux concentrations thus much clearer. The absence of three-dimensional mean-field structures was surprising because three-dimensional mean-field calculations have shown that the mean magnetic field develops structures along the direction of the imposed field (BKR). However, while the mean-field calculations have illustrated the nature of the instability, no systematic survey of solutions has yet been attempted. The purpose of this paper is therefore to clarify some still puzzling aspects concerning NEMPI. Note also that the large-scale flux concentrations observed in DNS of BKKMR have an amplitude of only 15\\,\\% of the local equipartition field. This implies that the flux concentrations we observe in DNS are often not strong enough to be noticeable without averaging.  In addition to the structures found in BKKMR, other types of structures have recently been reported in Large-Eddy Simulations (LES), which might also be an indication of NEMPI. We have here in mind the radiation magneto-convection simulations of Kitiashvili et al.\\ (2010), in which one sees the formation of whirlpool-like magnetic structures. Relevant to NEMPI is also the work of Tao et al.\\ (1998), who considered magneto-convection in the optically thick approximation and find a horizontal segregation into magnetized and non-magnetized regions. The size of the individual regions is such that they encompass several turbulent eddies. This phenomenon might therefore well be associated with an effect that could also be modelled in terms of mean-field theory. However, before we can make such an association, we need to find out more about the properties of NEMPI. In particular, we need to know what is the optimal magnetic field strength, what are the requirements or restrictions on the turbulent velocity, and, finally, how much density stratification is needed to make NEMPI work.  To connect the aforementioned requirements to DNS, we need to have a meaningful parameterization of the turbulence effects. The work done so far has been focussing on measuring a reduction of the turbulent pressure and effective mean magnetic pressure as a function of the local mean magnetic field strength. The shape of the resulting dependence of the effective mean magnetic pressure on the mean magnetic field has been matched to a specific fit formula that can be characterized by two fit parameters that, in turn, can be linked to the minimum effective mean magnetic pressure and critical field strength above which the effect is suppressed. However, there have been indications that this parameterization is not unique and that different combinations of the two fit parameters can result in similar values of minimum effective pressure and the critical field strength. The question therefore arises whether this apparent degeneracy affects the properties of NEMPI.  We mentioned already the fact that NEMPI is capable of exciting three-dimensional structures that show variation along the direction of the mean magnetic field. This would give rise to the worry that the two-dimensional results presented so far may not reflect the properties of the fastest growing mode and may therefore not be relevant to describing NEMPI. However, as will be discussed in this paper, this is not the case, because the degree to which three-dimensional modes are excited depends on the sign and magnitude of one of the turbulence parameters, namely the term characterizing turbulence effects on the magnetic tension force, and that simulations indicate that this sign is not favorable for exciting three-dimensional modes (BKKR, K\\\"apyl\\\"a et al.\\ 2012). Before we begin addressing the various points, we discuss first the mean-field model and basic setup. ",
        "conclusions": "The present work has clarified a number of puzzling aspects of NEMPI. Firstly, it is now clear that we can proceed with two-dimensional mean-field simulations as long as we know that $\\qsz=0$ (or negative). However, this may not always be the case. The fact that three-dimensional structures can emerge from NEMPI was initially thought to be an interesting aspect, because it could readily explain the formation of bipolar regions (BKR). However, given that simulations now indicate that $\\qs\\approx0$ (or perhaps even negative), this proposal would thus not be an option, unless some other as yet unexplored effect begins to play a role. In principle, all turbulent transport processes are nonlocal and must be described by a convolution with the mean field rather than a multiplication (Brandenburg et al.\\ 2008). In Fourier space, the convolution corresponds to a multiplication with a wavenumber-dependent turbulent transport coefficient. Thus, the idea of explaining bipolar regions would again become viable if this effect only existed at small and intermediate length scales. Clarifying this would be a task for future simulations, because none of the currently available techniques are yet equipped to address this possibility.  Next, we have seen that the degeneracy in the fit formula used for $\\qp(\\beta)$ and $\\Peff(\\beta)$ is significant in that different combinations of $\\qpz$ and $\\betap$ result in similar values of $\\min(\\Peff)$ and $\\betacrit$, but the growth rates can still be quite different. This means that it is not sufficient to measure only $\\min(\\Peff)$ and $\\betacrit$. Instead, to characterize the functional form of $\\Peff(\\beta)$ more accurately, we need some other characteristics to represent the dependence of this function near $\\beta=0$. One such possibility is to use the field strength $\\betamin$ for which the minimum of the effective magnetic pressure is reached.  Knowing the value of $\\betamin$ has particular relevance in determining the height where NEMPI occurs. For a given value of the imposed field strength $B_0$, the condition $B_0/\\Beq(z_{\\cal P})=\\betamin$ determines the height $z_{\\cal P}$, where the effective magnetic pressure attains a minimum, and thus the height $z_B$, which tends to be 2--3 scale heights below $z_{\\min}$; see the lower panel of \\Fig{lamfit2b}. Therefore, the value of $B_0$ does not directly affect the growth rate of NEMPI.  Finally, we have tried to establish an approximate dispersion relation to estimate the growth rate of NEMPI as a function of turbulent viscosity, turbulent magnetic diffusivity, mean field strength, and the strength of stratification. This formula may serve as a first orientation and can hopefully be improved further with future simulations. This formula can also be useful in connection with analytic estimates concerning the regimes when NEMPI is expected in DNS or under other more realistic circumstances."
    },
    "1107/1107.0427_arXiv.txt": {
        "abstract": "{ A dedicated cosmic muon Monte-Carlo event generator CMSCGEN has been developed for the CMS experiment. The simulation relies on parameterisations of the muon energy and the incidence angle, based on measured and simulated data of the cosmic muon flux. The geometry and material density of the CMS infrastructure underground and surrounding geological layers are also taken into account. The event generator is integrated into the CMS detector simulation chain of the existing software framework. Cosmic muons can be generated on earth's surface as well as for the detector located 90~m underground. Many million cosmic muon events have been generated and compared to measured data, taken with the CMS detector at its nominal magnetic field of 3.8 T. }  \\FullConference{XXth Hadron Collider Physics Symposium \\\\ November 16 -- 20, 2009 \\\\ Evian, France}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5805_arXiv.txt": {
        "abstract": " ",
        "introduction": "In his famous sonnet, ``Bright Star'' (1819), John Keats addresses a star, lamenting of the transience of human emotions---and of human life itself---in contrast to the star's immutability:   \\begin{quote} Bright star, would I were steadfast as thou art---\\\\ Not in lone splendor hung aloft the night\\\\ And watching, with eternal lids apart,\\\\ Like nature's patient, sleepless Eremite\\ldots\\\\ \\ldots yet still steadfast, still unchangeable\\ldots \\end{quote}  \\noindent Many decades later, Robert Frost alluded to ``Keats' Eremite'' in his poem, ``Choose Something Like a Star'' (1947).  The poet queries a star (``the fairest one in sight''), pleading for a celestial lesson that ``we can learn/By heart and when alone repeat.''  He finds,  \\begin{quote} It gives us strangely little aid,\\\\ But does tell something in the end\\ldots\\\\ It asks of us a certain height,\\\\ So when at times the mob is swayed\\\\ To carry praise or blame too far,\\\\ We may choose something like a star\\\\ To stay our minds on and be staid. \\end{quote}  Both poets invoke a millenia-long, cross-cultural tradition of finding in the ``fixed stars'' a symbol of constancy; sometimes cold, sometimes comforting. But these poems of Keats and Frost bookmark a period of enormous change in our understanding of cosmic variability.  Already by Keats's time---marked by the discovery of invisible infrared and ultraviolet light in the Sun's spectrum (by Herschel, 1800, and Ritter, 1801), and by the dawn of stellar spectroscopy (Fraunhofer, 1823)---astronomers were discovering that there was quite literally ``more than meets the eye'' in starlight.  Later in the 19th century, long-exposure astrophotography extended the reach of telescopes and spectroscopes to ever dimmer and farther objects, and provided reproducibly precise records that enabled tracking of properties over time.  In the 20th century, advances in optics and new detector technologies extended astronomers' ``vision'' to wavelengths and frequencies much further beyond the narrow range accessible to the retina.  By mid-century, some of these tools became capable of short-time-scale measurements.  Simultaneous with these technological developments were theoretical insights, most importantly from nuclear physics, that unveiled the processes powering stars, processes with finite lifetimes, predicting stellar evolution and death, including the formation of compact, dense stellar remnants.  By the time of Frost's death (1963), astronomers had come to see stars as ever-changing things, not only on the inhumanly long billion-year time scales of stellar evolution, but even over humanly accessible periods of years, months, and days.  Within just a decade of Frost's death, the discoveries of pulsars, X-ray transients, and gamma-ray bursts revealed that solar-mass-scale objects were capable of pulsing or flashing on timescales as small as {\\em milliseconds}.  We now know that the ``fixed'' stars visible to the naked eye represent a highly biased cross-sectional sample of an evolving population of great heterogeneity.  The more complete astronomical census made possible by modern astronomical instrumentation reveals the heavens to be as much a place of dramatic---sometimes violent---change as a harbor of steady luminance. The same instrumentation also reveals subtle but significant change even among some of the visible ``fixed'' stars.  Here I will point a Bayesian statistical telescope of sorts at one particular area of modern time-domain astronomy:  periodic variability.  Even this small area encompasses a huge range of phenomena, as is the case in other disciplines studying periodic time series.  I will focus on two small but prominent corners of periodic astronomy:  studies of {\\em pulsars} (rapidly rotating neutron stars) and of {\\em extrasolar planets} (``exoplanets,'' planetary bodies revolving around other suns).  New and upcoming instrumentation are producing rich data sets and challenging statistical inference problems in both pulsar and exoplanet astronomy.  Bayesian methods are well-suited to maximizing the science extracted from the exciting new data.  The best-known and most influential statistical methods for detecting and characterizing periodic signals in astronomy use {\\em periodograms}.  In the next section I will take a brief, Bayesian look at periodograms; they shed light on important issues common to many periodic time series problems, such as strong multimodality in likelihood functions and posterior densities.  In \\S~3 I describe detection and measurement of pulsars using data that report precise arrival times of individual light quanta (photons), including Bayesian approaches to arrival time series analysis using parametric and semiparametric inhomogeneous Poisson point process models.  In \\S~4 I turn to exoplanets, where the most productive detection methods as of this writing find planets that are too dim to see directly by looking for the reflex motion ``wobble'' of their host stars. Here the data are irregularly sampled time series with additive noise, with very accurate but highly nonlinear parametric models for the underlying orbital motion.  I will briefly describe some key inference problems (e.g., planet detection and orbit fitting), but I will focus on application of Bayesian experimental design to the problem of adaptive scheduling of the costly observations of these systems.  A running theme is devising Bayesian counterparts to well-known frequentist methods, and then using the Bayesian framework to add new capability not so readily achieved with conventional approaches.  A final section offers some closing perspectives. ",
        "conclusions": ""
    },
    "1107/1107.2939_arXiv.txt": {
        "abstract": "We present an approach to building interferometric telescopes using ideas of quantum information.  Current optical interferometers have limited baseline lengths, and thus limited resolution, because of noise and loss of signal due to the transmission of photons between the telescopes.  The technology of quantum repeaters has the potential to eliminate this limit, allowing in principle interferometers with arbitrarily long baselines. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5587.txt": {
        "abstract": "In order to explore the properties of extreme outer stellar disks, we obtained ultra-deep $V$ and {\\it GALEX} UV images of 4 dwarf irregular galaxies and one Blue Compact Dwarf galaxy and ultra-deep $B$ images of 3 of these. Our $V$-band surface photometry extends to 29.5 magnitudes arcsec$^{-2}$. We convert the FUV and $V$-band photometry, along with \\ha\\ photometry obtained in a larger survey, into radial star formation rate profiles that are sensitive to timescales from 10 Myrs to the lifetime of the galaxy. We also obtained \\HI-line emission data and compare the stellar distributions, surface brightness profiles, and star formation rate profiles to \\HI-line emission maps, gas surface density profiles, and gas kinematics. Our data lead us to two general observations: First, the exponential disks in these irregular galaxies are extraordinarily regular. We observe that the stellar disks continue to decline exponentially as far as our measurements extend. In spite of lumpiness in the distribution of young stars and \\HI distributions and kinematics that have significant unordered motions, sporadic processes that have built the disks---star formation, radial movement of stars, and perhaps even perturbations from the outside---have, nevertheless, conspired to produce standard disk profiles. Second, there is a remarkable continuity of star formation throughout these disks over time. In four out of five of our galaxies the star formation rate in the outer disk measured from the FUV tracks that determined from the $V$-band, to within factors of 5, requiring star formation at a fairly steady rate over the galaxy's lifetime. Yet, the \\HI surface density profiles generally decline with radius more shallowly than the stellar light, and the gas is marginally gravitationally stable against collapse into clouds. Outer stellar disks are challenging our concepts of star formation and disk growth and provide a critical environment in which to understand processes that mold galaxy disks. ",
        "introduction": "\\label{sec-intro}  Outer parts of galaxy disks present an extreme environment for star formation. Dwarf irregular (dIm) galaxies already challenge models of star formation because of their low gas densities even in the central regions (Hunter \\& Plummer 1996; Meurer \\et\\ 1996; van Zee \\et\\ 1997; Hunter, Elmegreen, \\& Baker 1998). Outer parts of dIm galaxies with their exceptionally low gas densities, therefore, present a particularly difficult test of our understanding of the cloud/star formation process. Furthermore, knowledge of the star formation histories of outer stellar disks provides information on the way in which the disk has grown with time and tests current models for disk growth (see, for example, Azzollini \\et\\ 2008).  Stars have formed even to very low disk surface brightness levels. For example, Saha \\et\\ (2010) have detected an exponential stellar disk in the LMC that extends to 12 disk scale lengths, the equivalent of an $I$-band surface brightness of 34 mag arcsec$^{-2}$. Although the number of young stars relative to older stars declines with radius, stars as young as a few Gyr are found to about 8.3 disk scale lengths. %LMC scale length is 1.32 deg = 1.15 kpc. In the late-type spiral M33, Barker \\et\\ (2011) have found stars younger than 0.5 Gyr at 4 and 5 disk scale lengths, although the bulk of the stars are older. %M33 S1 50\\% 2-4 Gyr ago but younger stars are present. Stars younger than 0.5 Gyr in both. 4R_D %S2 90\\% have ages >2.5 Gyr. 5R_D And in M31, Richardson \\et\\ (2008) find evidence for star formation over the past 0.25-2 Gyr as far out as 56 kpc. Furthermore, satellite ultraviolet imaging has revealed extended ultraviolet disks in many spiral galaxies that indicate on-going star formation in outer disks (Thilker \\et\\ 2007). For example, {\\it GALEX} images of M83 revealed UV-bright stellar complexes to 7 kpc, about 4 times the radius at which most of the \\HII\\ regions were detected (Thilker \\et\\ 2005).  Given the importance of outer stellar disks to our understanding of galaxy formation and evolution, we have conducted a program to detect and measure the extreme outer stellar disks of a sample of dIm galaxies through ultra-deep imaging in the optical and ultraviolet. In addition we have mapped the galaxies in \\HI-line emission. We use these data to examine the structure, stellar populations, and star formation activity in extreme outer dIm disks.  Our program objects were selected from a large survey of dIm, Blue Compact Dwarf (BCD), and Sm galaxies (Hunter \\& Elmegreen 2006). They were chosen to be representative of the larger dIm sample in terms of absolute magnitude $M_V$ and star formation rate. They were also chosen to be relatively free of nearby bright stars that would compromise photometry at ultra-low light levels. The new deep imaging data will be referred to here as the ``deep'' data, while the shallower survey data of the larger sample will be referred to as ``survey'' data. The galaxies and several important characteristics are listed in Table \\ref{tab-sample}.  \\section {Observations} \\label{sec-obs}  \\subsection{Optical images}  \\subsubsection{Data acquisition and basic reductions}  The deep optical imaging observations are summarized in Table \\ref{tab-obs} and $V$-band and \\ha\\ images are shown in Figures \\ref{fig-d53} through \\ref{fig-izw115}. The \\ha\\ images are part of the survey data (Hunter \\& Elmegreen 2004). Reductions were done using the {\\it Image Reduction and Analysis Facility} (IRAF\\footnote{IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under cooperative agreement with the National Science Foundation.}).  In May 2002 we %Phil Massey kindly obtained %for us deep $B$ and $V$ images of DDO 133 and $V$ images of IZw 115 with the Mosaic Imager on the Kitt Peak National Observatory (KPNO) 4-m Mayall telescope. The Mosaic Imager consists of 8 2k$\\times$4k SITe CCDs and gives a field of view of 36\\arcmin. Because our galaxies are small on the sky, they were centered on one of the CCDs although the entire mosaic of CCDs was reduced together. Dome flats and twilight sky flats were used to determine the pixel-to-pixel variations and scaled to remove the pupil ghost. The images were corrected for geometrical distortions, a background offset was removed, and the images were stacked, with cosmic ray removal, to produce a single final image in each filter. The Mosaic Imager images were calibrated using field images and Landolt (1992) standard stars observed with the Lowell Observatory 1.1 m telescope. Because IZw 115 was not observed in $B$, a ($B-V$)$_0$ of 0.4 was assumed to calibrate instrumental $V$ magnitudes. Uncertainties in the calibration are given in Table \\ref{tab-obs}.  We obtained optical images of DDO 53, DDO 86, and NGC 4163 with the Cassegrain Focus Imager on the KPNO 2.1 m telescope. The T2KA CCD was used for $V$ observations of DDO 53 and DDO 86, and T2KB was used for observations of DDO 53 in $B$ and of NGC 4163 in $V$ and $B$. The 2048$\\times$2048 pixel format yields a 10.4\\arcmin\\ field of view. Landolt (1992) standard stars were also observed for photometric calibration, and the rms of the calibrations are given in Table \\ref{tab-obs}. We used observations of an illuminated white spot hanging in the dome (dome flats) to remove pixel-to-pixel variations, and dark sky flats to remove large-scale structure across the CCD.  Our experience has been that the best method for flat-fielding varies from telescope to telescope. Discussion of this point can be found in Massey \\& Jacoby (1992). In some cases, such as at the KPNO 4-m, bright twilight flats can be used to remove both the pixel-to-pixel variations and the large scale gradients. But in other cases, such as the KPNO 2.1-m, twilight flats do not work well  and blank-sky flats work best. Therefore, for the KPNO 2.1-m data, we constructed a dark sky flat for each observing run. We obtained multiple images with the same exposure time and moon conditions as the galaxy observations. ``Blank sky'' targets were used, but there were still many stars in the images. The telescope was stepped between exposures and the series of images was co-added with an algorithm that removed the stars by comparing the images pixel by pixel and rejecting pixels that deviated more than 3$\\sigma$ from the median, where the $\\sigma$ is determined from the noise and gain parameters of the CCD. Because the sky flats showed no structure changes from night to night, a single sky flat was constructed for an entire observing run. The resulting sky flat was smoothed. We normalized the sky flat to a mean value of one near the center, and divided the object images by this sky flat.  In the case of the $V$-band image of DDO 53, however, we discovered after the observing run that the ``blank sky'' was contaminated by Milky Way nebular structure (4$^{\\rm h}$ 29$^{\\rm m}$ 45$^{\\rm s}$, 54\\arcdeg 15\\arcmin 33\\arcsec [J2000]). The Milky Way emission produced diffuse filamentary structure across the field of view and so could not be used to construct a dark sky flat. On the other hand, the galaxy itself was small and contained in the center of the field of view; even at low light levels the galaxy occupied only 14\\% of the pixels of the image. Therefore, for this galaxy/filter only, we used the galaxy images themselves to construct a dark-sky flat: we combined the images to produce a single star-free average and fit the non-galaxy portion of the image with an order 5 Legendre function, interpolating across the area of the image containing galaxy. The galaxy-free region was defined by looking at the extent of the galaxy with the image displayed to see faint noise variations in the background.  We observed the galaxies only at relatively low airmass ($\\leq1.4$), high transparency, and no moonlight. In all cases we took a series of exposures of the galaxy, and the telescope was stepped between exposures in order to move the galaxy around the chip and average over flat-fielding imperfections. We only retained images with less than 7\\% transparency differences and with point-spread-function FWHM less than 50\\% greater than the best in the series. The flat-fielded galaxy images were shifted to align the stars and co-added to build signal-to-noise. We used an algorithm to remove cosmic rays and adjusted parameters so that the photometric properties, as determined in a simple average of the images, was preserved.  In spite of our great care in determining flat-field data, the night-sky brightness at KPNO (or any dark site) is about 22.5 mag arcsec$^{-2}$ (see, for example, Neugent \\& Massey), and so achieving 29.5 mag arcsec$^{-2}$ requires going 7 magnitudes below the sky.  No flat-field will be good to 1 part in 6000 (10\\% photometry) at these light levels. Fortunately, it does not have to be.  The key to measuring these low light levels is not in the flat-fielding, but rather in the sky subtraction.  The flat-fielding helps in getting the sky relatively flat from one side of the frame to the other, but even if there are 1\\% residual gradients in the flat-fielding, this is dealt with by fitting the sky with a 2-dimensional function, as we did.  Therefore, we edited foreground stars and background galaxies from the images and determined a two-dimensional fit of the sky pixels in the edited image---interpolating across the galaxy. This procedure works because the galaxies are a small portion of the total field of view of the image: approximately 14\\% for DDO 53, 12\\% for DDO 86, 10\\% for DDO 133, 26\\% for NGC 4163, and 4\\% for IZw 115 and its detached region. It is necessary because there is structure in the sky. Typical variations of the background across the image were of order 0.2-0.6\\%. We used a Legendre function with the minimum order that produced a satisfactory sky. Typical orders were 3-5. We determined the success of a fit to the sky empirically by examining row and column cuts through the sky-subtracted image. We expected the background to be flat with noise around zero, and we varied the x and y orders until we achieved that. This procedure produced an image of the sky. We then subtracted the sky image from the original co-added image and from the edited image to yield sky-subtracted images with and without the foreground/background object contamination.  A legitimate question is whether or not our fits to the sky structure are affected by the presence of the galaxy and/or contaminating stars in the field.  That can be answered best by comparing results obtained on the same galaxy using different telescopes/instruments with differing fields of view. This comparison is made in \\S \\ref{sec-optsurfphot}. We further gain confidence in our results in the outer disks of the dwarfs from the fact that we see similar profiles from data obtained with different telescopes/instruments and reduced with different types of sky flats (DDO 133 and IZw 115 imaged with the Mosaic Camera on the KPNO 4 m and reduced with twilight flats, and the other galaxies imaged with the Cassegrain Focus Imager on the KPNO 2.1 m telescope and reduced with dark sky flats).  \\subsubsection{Surface photometry} \\label{sec-optsurfphot}  Removal of background galaxies and foreground stars is essential for meaningful azimuthally-averaged surface photometry. We attacked this problem in two ways for each image. First, we used the edited sky-subtracted images, with some additional masking of small areas in close to the galaxy. A problem with this approach is in distinguishing, without radial velocity information, stellar clusters in the galaxy from foreground/background object contamination. We used the various passband images of the galaxy at our disposal to make the best guess as to whether a discrete object in the image was a part of the galaxy or foreground/background contamination. The area surrounding the galaxy is a good clue to the type of contamination we should expect across the galaxy. Nevertheless, subjective judgement was involved in this step.  The second method we used to account for foreground/background contamination was to use the unedited sky-subtracted image and determined an average ``sky'' value around the galaxy that included the foreground/background component. This statistically corrects for interlopers, and assumes that the fields around the galaxy represent the contamination across the galaxy itself. The problem with this method is that a given annulus in the surface photometry might not contain an exactly representative amount of light from extraneous objects, and a single foreground star or background object in the faint, outer parts of the galaxies can severely affect the surface photometry. For this reason, we prefer the surface photometry obtained with the first method of editing contamination from the images, but we applied both methods. In fact the two methods produced similar results. The comparison is shown in Figures \\ref{fig-comparemethods} and \\ref{fig-compareuvmethods} for the $V$-band and {\\it GALEX} NUV photometry, respectively. The surface photometry is cut off where there is a large jump in the uncertainty, usually around 1 magnitude arcsec$^{-2}$.  Because dIm galaxies are lumpy, we determined the center of the galaxy, position angle (P.A.), and ellipticity from an outer isophote in a contour plot of a smoothed version of the $V$ image. Then we integrated the photometry in fixed ellipses that increase in semi-major axis length. The step size is given in Table \\ref{tab-obs}, and was chosen to be $\\geq$150 pc and, as seen by eye, to cover a sufficient area of the galaxy to average over the most severe local fluctuations. We geometrically matched the continuum-subtracted \\ha\\ images to the $V$-band images and applied the same ellipse integrations. The outer most ellipse is shown superposed on the $V$, NUV, and \\ha\\ images in Figures \\ref{fig-d53} through \\ref{fig-izw115}.  We corrected the photometry for reddening using the foreground reddening, given in Table \\ref{tab-sample}, plus 0.05 mag in E($B-V$) for internal reddening in $B$ and $V$ and 0.1 mag in E($B-V$) for internal reddening in \\ha. This  level of internal reddening is consistent with measurements of the Balmer decrement in \\HII\\ regions in a sample of 39 dIm galaxies (Hunter \\& Hoffman 1999). There, the average reddening in \\HII\\ regions is 0.1, and we have taken half this to represent the stars outside of \\HII\\  regions. We used the Cardelli \\et\\ (1989) extinction curve.  We used a single extinction correction for all radii within each galaxy. This is justified by the lack of metallicity gradients in dwarfs (see, for example, Croxall \\et\\ 2009) and the correlation between metallicity and attenuation seen in spirals (Boissier \\et\\ 2007). On the other hand, little is known about the outer disks of dwarfs, beyond where \\ha\\ emission is no longer detected, although there are some hints of dust emission even there (see, for example, Hinz \\et\\ 2006). Boissier \\et, based on {\\it IRAS} observations, suggest that there is no attenuation in the outer disks of two dwarfs (IC1613, WLM), but uncertainties are large due to low {\\it IRAS} fluxes.  There is a region of stellar emission near, but detached from, IZw 115 that we did not include in the galactic photometry. It is located at 15$^h$32$^m$51.0$^s$ 46\\arcdeg25\\arcmin51\\arcsec\\ (2000), 101.5\\arcsec\\ ($=$6.0 kpc) from the center of IZw 115, southwest along the galaxy's major axis, and is marked in images in Figure \\ref{fig-izw115}. In a circular aperture of radius 14.85\\arcsec, the region has a $V$ magnitude of 20.39$\\pm$0.04 and a surface brightness $\\mu_V$ of 26.9 mag of one arcsec$^2$. If it is at the distance of IZw 115, it has an $M_V$ of $-10$ and a diameter of several kpc. Thus, it could be a very small dwarf companion.  DDO 133 is clearly barred, and the outer isophote that was used to determine the galaxy morphology is located beyond the bar. The bar has a P.A. that is $-13$\\arcdeg, an 11\\arcdeg\\ twist from the outer isophote. The center of the bar is located at 12$^h$32$^m$54.5$^s$ 31\\arcdeg32\\arcmin27\\arcsec, 18.8\\arcsec\\ ($=$560 pc) from the center of the outer isophote. The minor-to-major axis ratio $b/a$ of the bar is 0.44, compared to 0.69 for the outer isophote, so the outer isophote is rounder. The major axis of the bar is 133\\arcsec\\ ($=$3.9 kpc) long. We experimented with changing the centering and P.A. of the photometry ellipses, but found that there was no significant change in the photometry for a reasonable alternate choice of the center (10\\arcsec\\ offset to the southwest) and P.A. (5\\arcdeg\\ counter-clockwise).  DDO 86 presents a special challenge to determining the position angle of the major axis. At an outer isophote ($\\mu_V\\sim27.2$ mag arcsec$^{-2}$), the P.A. of the major axis is 85\\arcdeg\\ and the minor-to-major axis ratio $b/a$ is 0.86. This is close to the galaxy morphology determined in the survey data: P.A.$\\sim$72\\arcdeg\\ and $b/a\\sim$0.82 (Hunter \\& Elmegreen 2006). However, in the deeper imaging data that we have here, we see that there is a faint extension of starlight to the north. At an isophote level that includes this extension ($\\mu_V\\sim28.4$ mag arcsec$^{-2}$), the P.A. of the galaxy rotates clockwise to 32.5\\arcdeg, although the shape remains the same ($b/a\\sim0.87$). An inner isophote ($\\mu_V\\sim25.6$ mag arcsec$^{-2}$) has a major axis that is rotated counter-clockwise from both of these (113.5\\arcdeg) and the galaxy is more elongated ($b/a\\sim0.65$). Finally, the kinematical major axis determined from the \\HI\\ velocity field has a P.A. of $-46$\\arcdeg. It is possible that this sequence of optical P.A.s---rotating clockwise from higher surface brightness to fainter---and change in shape---becoming rounder in the middle and outer parts---represents an offset bar; the center of the inner isophote is located 360 pc to the southwest of the center of the extreme outer isophote. For the surface photometry we have chosen to use the P.A. from our \\HI\\ kinematic analysis.  Since we have surface photometry from the survey images of each of these galaxies, we compare the deep imaging photometry with that from the survey data in Figure \\ref{fig-comparemethods}. We see that the agreement is quite good in most cases. In addition, the breaks in the surface brightness profiles seen in the survey data of DDO 86 and DDO 133 are there at the same radius and surface brightness level in the deep data. And, the azimuthally-averaged $B-V$ colors in DDO 133 match closely the survey data in the region of overlap. The exception to the good agreement is DDO 53. Beginning at a radius of 1.2\\arcmin\\ (1.25 kpc) and a $\\mu_V$ of 25.8 mag arcsec$^{-2}$ the new surface photometry is systematically fainter. The differences are greater than the uncertainties in the photometry. The survey $B-V$ color is also systematically redder than the deep photometry beginning at a radius of 0.77\\arcmin\\ (0.81 kpc; $\\mu_V\\sim24.9$ mag arcsec$^{-2}$). In the survey data, there is a gradient in $B-V$, with the galaxy becoming redder with radius. In the deep data, the color is relatively constant. We note that the integrated $B-V$ of DDO 53 from the deep image (0.30$\\pm$0.01, not corrected for reddening) is close to that given by de Vaucouleurs \\et\\ (1991; 0.35), while that from the survey data (0.48) is significantly redder. Therefore, we suspect that the deep photometry is better than that measured from the survey data for this galaxy in the region of overlap.  From the ellipse photometry we can derive several quantities. First, the last ellipse yields integrated photometry and that is given in Table \\ref{tab-phot}. Second, fits to the surface brightness profiles yield disk scale lengths $R_D$ and central surface brightnesses $\\mu_{V,0}$. These are listed in Table \\ref{tab-disk}. Four of the 5 galaxies---DDO 53, DDO 86, DDO 133, and NGC 4163---exhibit double exponentials. The profile in the outer disk falls off more steeply than that in the inner disk in DDO 53, DDO 86, and DDO 133. In NGC 4163 the outer profile is shallower than the inner. DDO 86 and DDO 133 were already known to have double exponentials from the survey data, but the survey data of DDO 53 did not go out far enough to give a hint of the break in the profile. The break in NGC 4163 is more subtle than in the other three galaxies. Nevertheless, the scale lengths are different---$0.26\\pm0.02$ kpc versus $0.31\\pm0.004$ kpc, and the central surface brightnesses implied by each differ by 0.54 magnitude. For those four galaxies with breaks, the $R_D$ and $\\mu_{V,0}$ are given for both the inner and outer parts of the profiles, along with the break radius $R_{Br}$ at which the slope changes. In addition DDO 53 and IZw 115 show breaks at small radii, a little over one disk-scale length that coincide with concentrations of star formation. These are marked in the surface brightness plots, but not included in Table \\ref{tab-disk}. The surface photometry and color profiles are shown in Figures \\ref{fig-d53sb} through \\ref{fig-izw115sb}.  \\subsection{Ultraviolet images}  We obtained deep images of our sample galaxies with the {\\it Galaxy Evolution Explorer} satellite ({\\it GALEX}; Martin \\et\\ 2005) to trace and characterize star formation into the outer disks of dwarf galaxies where \\ha\\ may not be an effective tracer of recent star formation (see, for example, Thilker \\et\\ 2005, 2007a; Boissier \\et\\ 2007). The UV---ultraviolet light coming directly from OB stars---can easily trace young stars in the outer galaxy, and the UV has the advantage in a patchy star-forming environment of integrating over a longer timescale than does \\ha\\ (10 Myrs).  The galaxies are listed in Table \\ref{tab-galex} along with the exposure times and tile name of the {\\it GALEX} images. {\\it GALEX} imaged simultaneously in two channels: FUV with a bandpass of 1350--1750 \\AA, an effective wavelength of 1516 \\AA, and a resolution of 4.0\\arcsec\\ and NUV with a bandpass of 1750--2800 \\AA, an effective wavelength of 2267 \\AA, and a resolution of 5.6\\arcsec. The images were processed through the {\\it GALEX} GR4/5 pipeline and were retrieved as final intensity maps with a 1.5\\arcsec\\ pixel scale. The {\\it GALEX} field of view is a circle with 1.2\\arcdeg\\ diameter, and we have extracted a portion around our target galaxies.  We corrected the UV photometry for extinction using the same total reddening E($B-V$)$_t$ as was used for correcting the optical. We combined the E($B-V$)$_t$ with the extinction law of Cardelli \\et\\ (1989) to produce the extinction $A_{FUV}=8.24{\\rm E}(B-V)_t$, and, interpolating over the 2175 \\AA\\ bump, $A_{NUV}=7.39{\\rm E}(B-V)_t$. However, the FUV extinction law is known to vary from galaxy to galaxy, and the NUV filter straddles the 2175 \\AA\\ extinction feature, which also varies not only from galaxy to galaxy but from place to place within galaxies (for example, Gordon \\et\\ 2003). Star formation activity appears to play as important a role, perhaps even more important, than metallicity in determining the UV extinction curve and the strength of the 2175 \\AA\\ bump. In the metal-poor SMC, for example, Gordon \\& Clayton (1998) have shown that lines of sight through the actively star forming bar have UV extinction curves that are linear with 1/$\\lambda$, rising more steeply into the FUV than in the Milky Way, and no 2175 \\AA\\ bump is present, while a sight line through the quiescent wing has an extinction curve that is similar to that in the Milky Way with a shallower slope into the FUV and a 2175 \\AA\\ bump. In our dwarf galaxies, the level of star formation activity is highly variable from galaxy to galaxy and from place to place within a galaxy, potentially resulting in a variable FUV extinction law. Wyder \\et\\ (2007) have dealt with this for a large sample of galaxies by adopting the Cardelli \\et\\ extinction law for the Milky Way and determining the best average reddening from convolving theoretical galaxy spectral energy distributions (SEDs) with the {\\it GALEX} filter transmission curves. In Hunter, Elmegreen, \\& Ludka (2010), we examined the results of convolving different SMC, LMC, and Milky Way extinction curves (Cardelli \\et 1989, Gordon  \\et\\ 2003) with constant star formation rate galaxy SEDs constructed from the Bruzual \\& Charlot (2003) stellar population library and with the {\\it GALEX} filter transmission curves. The $A_\\lambda/{\\rm E}(B-V)$ varies by 1.5 magnitudes depending on the extinction curve. Considering the variety of environments represented by our sample and the fact that extinction is minimal except for a few galaxies with higher than average foreground extinction, we decided to adopt the Wyder \\et\\ extinctions: $A_{FUV}=8.24{\\rm E}(B-V)_t$ and $A_{NUV}=8.2{\\rm E}(B-V)_t$. For most of our galaxies, the variations represented by alternate extinction laws result in differences of order 0.05 mag.  Surface photometry of the NUV and FUV images proceeded in the same manner as for the optical images. We edited foreground/background objects from the images, fit the cleared area around the galaxy, and subtracted this from the edited and from the unedited images. For the unedited sky-subtracted image, we determined the average background light from contaminating objects and subtracted this from the image. We then geometrically transformed the UV images to match the $V$-band scale and orientation and measured surface photometry on the edited and unedited images using the same ellipse parameters as for the $V$-band images. The two methods of accounting for background/foreground contamination are shown in Figure \\ref{fig-compareuvmethods}. The results of the two methods agree well. Integrated photometry are given in Table \\ref{tab-phot}. The surface photometry and color profiles are shown in Figures \\ref{fig-d53sb} through \\ref{fig-izw115sb}.  Since we are looking at faint light levels in outer disks, we need to consider whether the UV light we see truly comes from young stars or whether there might be other sources of light that dominate out there. One such potential source is white dwarf stars. DDO 133, as a test case, has an $FUV-NUV$ color of order 0.1 mag in the outer disk, not counting the colors with relatively high uncertainties. A white dwarf of that color has an effective temperature of order 16,000 K. At a distance of 10 pc, a star with $\\log g=6$ and a radius $R=R$\\solar\\ would have an absolute NUV AB magnitude of $M_{NUV}=1.84$ (L.\\ Bianchi, private communication). Assuming a mass for the white dwarf of 0.7 $M$\\solar, the star should have a radius of 0.1$R$\\solar. Thus, the apparent NUV magnitude would be 35.7. To reproduce a $\\mu_{NUV}$ of 28 mag arcsec$^{-2}$ at  a radius of 2$R_D$ in DDO 133, we would need about 1.4 white dwarfs pc$^{-2}$. That is 2 times higher than the density of white dwarfs in the solar neighborhood. For white dwarfs with smaller radii or higher $\\log g$, the number of white dwarfs that are needed goes up, so this is probably a lower limit. However, the solar neighborhood should have a much higher stellar density than the outer disk of a dwarf. In the solar neighborhood $\\mu_V$ is 22.7$\\pm$0.2 (Ishida \\& Mikami 1982), and that is a factor of 30 higher than the surface brightness in DDO 133 at the radius we are considering ($\\mu_V=26.5$). So, the stellar density of white dwarfs needed to account for the UV light in DDO 133 fails by a factor of about 60. Even 3 magnitudes fainter---at the edge of our detection, the numbers needed are 4 times too high and so it is unlikely that white dwarfs alone could produce the UV starlight that we see. We also considered horizontal branch stars since these dominate the UV light in, for example, M32 (Brown \\et\\ 2000). However, the $FUV-NUV$ color of these stars are of order 2 and redder. This color is far redder than what we see in dwarfs.  \\subsection{\\HI}  Three of the galaxies in our sample---DDO 53, DDO 133, NGC 4163---are part of the LITTLE THINGS ({\\it LITTLE: Local Irregulars That Trace Luminosity Extremes; THINGS: The HI Nearby Galaxy Survey}) project, a multi-wavelength survey of dwarf galaxies (LITTLE THINGS: Hunter \\et, in preparation; THINGS: Walter \\et\\ 2008). LITTLE THINGS has \\HI\\ observations at additional VLA arrays, including B array, that provide higher spatial resolution. They will provide much better maps for analyzing the relationship between the gas and star formation. Here we restrict ourselves to the data obtained for this project and examine general correlations between the stars and the gas.  \\subsubsection{Calibration and Mapping}  We obtained Very Large Array (VLA\\footnote{ The VLA is a facility of the National Radio Astronomy Observatory (NRAO). The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}) observations of \\HI-line emission in our galaxies. The dates and configurations of the L-band observations are given in Table \\ref{tab-vla}. The total bandwidth was 1.56 MHz, and 2 IFs were used. The data were on-line Hanning smoothed, and the channel separation was 12.2 kHz (2.6 \\kms). The data from 2006 and 2008 were obtained when the VLA was in transition to the EVLA. Unfortunately, in converting the signal from the EVLA antennas, power was aliased into the lower 500 kHz of the bandwidth. The consequence of this is that all EVLA-EVLA baselines had to be flagged and not used (see Ficut-Vicas et al., in preparation, for a discussion of how to deal with this problem).  We subtracted the continuum emission in the {\\it uv}-plane using line-free channels on either end of the spectrum and, in the case of DDO 53, combined the data from the two arrays. To make maps, we employed a routine in NRAO's Astronomical Image Processing System (AIPS) that enables one to choose a sample weighting that is in between the standard ``natural'' weighting, which gives the highest signal-to-noise but at the expense of long wings in the beam, and ``uniform'' weighting, which gives the best resolution but at the cost of signal-to-noise and the presence of negative sidelobes. We chose a sample weighting that gives a formal increase of only 5\\% in noise yet with a significant improvement in beam profile over what would have been achieved with ``natural'' weighting.  The resulting synthesized beam FWHM are given in Table \\ref{tab-vla}. We deconvolved the maps until there were roughly comparable numbers of positive and negative components.  To remove the portions of each channel map without emission that are only contributing noise to the map, we used the maps smoothed to twice the beamsize for a conditional transfer of data in the unsmoothed maps. The channel maps were blanked wherever the flux in the smoothed maps fell below 2.5$\\sigma$. Flux-weighted moment maps were made from the resulting data cube. Integrated emission---moment zero---maps are shown in Figures \\ref{fig-d53vf} through \\ref{fig-izw115vf}, along with velocity field (moment one) and velocity dispersion (moment two) maps. Integrated \\HI\\ contours superposed on the $V$-band images are shown in Figures \\ref{fig-d53} through \\ref{fig-izw115}.  \\subsubsection{Velocity fields} \\label{sec-vf}  Position velocity diagrams (panel e), velocity fields (panels b and d), and velocity dispersions (panel c) are shown in Figures \\ref{fig-d53vf} to \\ref{fig-izw115vf}. Panel (b) shows the standard intensity-weighted mean velocity field, and panel (d) shows the velocity field that results from fitting profiles with a Gauss-Hermite $h_3$ polynomial.  Because of the possible presence of multiple velocity components, we have deconvolved ordered from non-ordered motions using the method described in Oh \\et\\ (2008). First, an approximate rotation curve is derived from the major axis position-velocity diagram and this is used to create an artificial velocity field. Then in an automated method, we fit single Gaussians to the major velocity component in each profile across the galaxy. The velocities of these components are compared to the artificial velocity field. If they deviate more than a prescribed amount, the program examines the possibility that the bulk motion at that point is represented by a secondary velocity component. From this, we extract the bulk velocity field that then becomes the input velocity field for the next iteration, which is now fit using the ``tilted ring model'' rather than the major axis position-velocity diagram. This loop continues until the mean difference between successive velocity fields is less than three times the channel width. This gives a final bulk velocity field. Parameters derived from this process (center position, P.A., inclinations) are summarized in Table \\ref{tab-himaps}. Following Bureau \\& Carignan (2002) we corrected the rotation velocities for asymmetric drift, and the final rotation curves are shown in the graphs in Figures \\ref{fig-d53sb} through \\ref{fig-izw115sb}.  The velocity fields of DDO 86 and DDO 133 exhibit well-behaved ordered rotation. In IZw 115 the velocity field is less well-behaved, but still ordered motion is present. The velocity fields in DDO 53 and NC 4163, on the other hand, are quite peculiar. Here we discuss each galaxy individually:  {\\it DDO 53}. DDO 53 is a part of the THINGS study of \\HI\\ in nearby galaxies (Walter 2008) and has been analyzed by Oh \\et\\ (in preparation). The kinematical axis of DDO 53 is at a position angle of 155\\arcdeg, while the optical major axis is at an angle of 81\\arcdeg, a difference of 74\\arcdeg. Furthermore, the fit to the \\HI\\ velocity field yields a variable inclination angle. The kinematic center and the optical center are 32\\arcsec\\ apart, with the kinematic center being NE of the optical center. The fit to the velocity field is quite poor, and the rotation curve is only meaningful out to a radius of 0.9 kpc.  {\\it DDO 86}. The bulk velocity field of DDO 86 is well determined with clean fits to the position angle of the kinematic axis, the center position, and the inclination of the \\HI\\ disk. The optical, however, is more problematic, as discussed in \\S \\ref{sec-optsurfphot}: The $V$-band major axis rotates clockwise from higher surface brightness to fainter. This may indicate the presence of an off-center bar. So, for the optical and UV surface photometry we used the P.A.\\  derived from the \\HI\\ kinematics. The optical center determined from outer isophotes agrees within 14\\arcsec\\ with the \\HI\\ center, with the \\HI\\ center being NW of the optical.  {\\it DDO 133}. The bulk velocity field of DDO 133 is well-behaved. The kinematic P.A., inclination, and center agree very well with those determined for the optical.  {\\it NGC 4163}. Although rotation is seen in the gas, the velocity field is so peculiar that parameters for the tilted ring model---center, P.A., and inclination--were taken from the optical.  {\\it IZw 115}. IZw 115 shows rotation, although the velocity field is rather peculiar. The position angle of the kinematic axis was found to vary between 50\\arcdeg\\ and 78\\arcdeg, while the optical major axis is at an angle of 41\\arcdeg. On the other hand, the inclinations agree within 7\\arcdeg, and the centers, within 4\\arcsec. The detached region seen in the optical and NUV (Figure \\ref{fig-izw115}) shows a hint of rotating independently of the main body of IZw 115. Panel (d) in Figure \\ref{fig-izw115vf}, which shows the velocity field resulting from fitting profiles with a Gauss-Hermite $h_3$ polynomial, can be interpreted as showing overall rotation with an axis at a position angle of roughly 90\\arcdeg, clearly distinct from IZw 115. This could indicate that the detached region is a small companion galaxy that is interacting with IZw 115.  \\subsubsection{Surface density profiles}  The azimuthally-averaged \\HI\\ surface densities $\\Sigma_{HI}$ and velocity dispersions $\\sigma_{vel}$ of the rotational component were determined using the center, P.A., and inclination angle given in Table \\ref{tab-himaps}. The $\\Sigma_{HI}$ was corrected for He %and, statistically, for H$_2$ to produce the gas surface density $\\Sigma_{gas}$. %The H$_2$ component is taken to be 18\\% of $\\Sigma_{HI}$, %from an average of observations of the SMC and LMC (Cohen \\et\\ 1988, Rubio \\et\\  1991). The gas surface density profiles and average velocity dispersions are shown as a function of radius in Figures \\ref{fig-d53sb} through \\ref{fig-izw115sb}. ",
        "conclusions": "\\subsection{Radial Profiles}  Deep $V$-band images of 4 dwarf irregular galaxies and one BCD galaxy have been used to measure surface photometry down to $\\sim29.5$ mag arcsec$^{-2}$.  Deep $B$-band images were also used for 3 of these galaxies. The surface brightness distributions in the $V$-band were found to be exponential in radius out to 5 inner disk scale lengths ($R_{\\rm D}$) in two cases (DDO 53, DDO 86), $8R_{\\rm D}$ in two other cases (NGC 4163, IZw 115), and $3R_{\\rm D}$ in one case (DDO 133). A clear break marks the transition from an inner exponential profile to an outer, steeper exponential profile in three cases. In two of these (DDO 53, DDO 86), the break occurs at $\\sim3R_{\\rm D}$, and in the third (DDO 133) it occurs at $\\sim1.8R_{\\rm D}$. In two cases where the $B-V$ color information goes beyond the break, one (DDO 53) has a transition in the color profile from a slight reddening in the inner disk (from $0.12\\pm0.005$ to $0.36\\pm0.03$, a difference of 0.24 mag) to an increasing blueness (down to $-0.13\\pm0.3$ at the last point shown, a difference of 0.49 mag) in the outer disk; the other (DDO 133) has a relatively flat color profile over all (variations up, down, up, and down again with radius covering an extreme of $0.47\\pm0.005$ to $0.30\\pm0.01$, a difference of 0.17 mag).  The remarkable regularity of these underlying exponential disks contrasts with the irregularity observed in this Hubble morphology class. Such regularity has been observed in the inner parts of dwarf galaxies for a long time (Patterson \\& Thuan 1996; Hunter \\et\\ 1998; Bremnes \\et\\ 1998, 1999, 2000;  Hopp 1999; van Zee 2000; Taylor \\et\\ 2005; Hunter \\& Elmegreen 2006; Amor\\'in \\et\\ 2007; but for exceptions see, for example, Hunter \\& Gallagher 1985; Hunter \\& Plummer 1996; Herrmann \\et, in preparation), but now we see that it extends into the extreme outer disk as well. Somehow the accumulated star formation and radial mass migration over a Hubble time converts the irregular activity at any one time into a standardized disk profile. This conversion seems to be done without grand design spiral or bar torques (only one case, DDO 133, is clearly barred and none have obvious spirals, but see below), and without significant shear in \\HI\\ velocity maps for at least the inner exponential parts. A lack of bars would mean that the Hohl (1971) mechanism for the formation of exponential profiles by bar and spiral torques is not operating. A lack of shear would suggest that the Lin \\& Pringle (1987) mechanism in which the star formation rate is proportional to the shear viscosity is also not operating.  {\\it GALEX} ultraviolet observations were also examined for these galaxies. In four cases, the NUV observations go out as far in the disk as the deep $V$-band images. The NUV profiles are also exponential and have about the same scale lengths as the $V$-band profiles. In two of the three cases with exponential breaks in $V$-band, there is also a break in the NUV profile at the same radius. In the one case with a $B-V$ color change at the break radius (DDO 53), there is no evident break in the NUV profile at all, even though it extends for one inner disk scale length beyond the $V$ break. This lack of a NUV break in DDO 53 could be related to the increasing blueness with radius seen in $B-V$: the blueness could compensate for the steeper mass profile seen in the $V$-band and make the NUV-band appear continuous.  The similarity between the NUV and $V$-band profiles introduces another puzzle about the formation and evolution of these galaxies. The average radial profile of the star formation rate over the last several tenths of a Gyr, as observed in the NUV, follows the underlying mass distribution established by star formation and radial migration over the last Gyr or more, as observed in $V$. There was no radial mixing of these disks on timescales longer than $\\sim1$ Gyr because that would significantly affect the color profiles. For example, the outer exponential disk should be redder than the inner disk if it was formed or augmented by spiral arm scattering of inner disk stars $\\sim1$ Gyr old or more (e.g., Ro\\v skar et al. 2008). Such a color change has been observed for spiral galaxies (Bakos, Trujillo \\& Pohlen 2008), but in DDO 133 there is no color change and in DDO 53 the outer disk is bluer. Neither galaxy has spiral arms anyway.  Consider an example using the population synthesis models of Bruzual \\& Charlot (2003). For a Salpeter IMF (the results are essentially the same for a Chabrier IMF), a burst of star formation older than $10^{8.5}$ yrs has $NUV-V>2$ mag, and one older than 1 Gyr has $NUV-V>3.5$ mag. Young stars have $NUV-V<-1$.  These values were determined by integrating the model stellar population spectra over the filter functions used in the two passbands.  A plot of them is shown in Figure 21 of our previous paper (Hunter et al. 2010). If we took a group of stars that formed $10^{9}$ years ago and moved it to a different place in the disk, then that group would be 3.5 magnitudes fainter in NUV than in V because of population fading.  This 3.5 mag difference means that the displaced group may still stand out in V but be nearly invisible in NUV. Then the two passbands will have different intensity distributions. This difference of 3.5 mag is so large that even with more continuous star formation than a burst, and with some mixing of the group into other stars, there should still be a significant difference between the NUV and V intensities of stars $>1$ Gyr old. Our observations do not show this difference, however. The similarity between the NUV and V profiles in our galaxies implies that there has not been much migration of stars from the inner disk to the outer disk on timescales longer than $\\sim1$ Gyr.  The observation of similar radial profiles in NUV and $V$-bands (and also in $J$-band for at least the inner regions -- see Hunter \\& Elmegreen 2006) suggests that star formation occurred at a steady rate. Leroy et al. (2008) also found that the star formation rate in dwarfs was about constant over a Hubble time. In only one of our cases (NGC 4163) is the average star formation rate observed in the NUV different from the average observed in the $V$-band by as much as a factor of 10 (the UV rate is less than the $V$ rate at mid-disk), but in this case, the overall kinematics of the \\HI\\ is very peculiar, suggestive of a major disruption in the recent past. Still, NGC 4163 has a near-perfect exponential disk in the $V$-band out to 8 scale lengths with no break; it is the NUV profile that deviates from exponential at mid-radius.  H$\\alpha$ observations are available for all of these galaxies from our previous surveys. The H$\\alpha$ radial profiles extend out to only $2R_{\\rm D}$ or $3R_{\\rm D}$. In the case with the shortest $V$-band break radius (DDO 133), the H$\\alpha$ extends beyond the breaks in $V$ and NUV and also has a break at the same radius. In all of the galaxies, the H$\\alpha$ profile is roughly similar to the NUV profile, including local deviations from the exponential. This can be seen in Figure \\ref{fig-sfrratfuvha} where we plot the ratio SFR$_{FUV}$/SFR$_{H\\alpha}$. The implication is that current star formation rates, seen in H$\\alpha$, are about the same at each radius as the average rate over the last few hundred Myr. The H$\\alpha$ rates today are also about the same, to within a factor of $\\sim3$, as the $V$-band long-term average rates in three of the galaxies (see Figure \\ref{fig-sfrrat}). The two galaxies where these rates differ the most are NGC 4163, which has highly disturbed kinematics as mentioned above (and a current rate in H$\\alpha$ and the UV that is much lower than the $V$-band average at mid-radius), and the BCD galaxy, IZw 115, which also has a low $H\\alpha$ compared to $V$ in the inner parts (although the UV rate is the same as the $V$ rate there). IZw 115 is evidently post-starburst.  \\subsection{Hydrogen Gas}  The final piece of evidence related to star formation and galactic structure is the gas distribution. For this, we examined \\HI\\ observations taken at the VLA for this project. The \\HI\\ radial profiles are smoothly declining but flatter than the stellar light profiles in DDO 86, NGC 4163 and IZw 115, and about the same as the stellar light in DDO 133. In DDO 86, there is an \\HI\\ hole in the center and a ring around it that is not present in the optical. These trends show up in the SFR-$\\Sigma_{\\rm gas}$ diagrams of Figure \\ref{fig-nhisfr} too.  DDO 86, DDO 133, and IZw 115 have steep near-power-law relations between the $V$-band or FUV SFRs and the \\HI\\ column densities, with a slopes of $3\\pm1$ to $4\\pm1$. NGC 4163 and IZw 115 have a scattering of points with a somewhat more organized relation for FUV than $V$-band SFRs. The near-power-laws are consistent with other studies of star formation and \\HI\\ column density, which also show steep slopes. Bigiel et al. (2010) got a slope of 1.7 for the relation between SFR per unit area and \\HI\\ column density in faint outer regions. Their plot of $\\Sigma_{HI}$ versus $\\Sigma_{SFR}$ for dwarfs in their Figure 8 is similar in shape to ours for the FUV SFR in Figure \\ref{fig-nhisfr}. Star formation follows unseen CO (Bigiel et al. 2011), which is a small fraction of the \\HI\\ mass in these faint regions (Schruba et al. 2011). The timescale for conversion of \\HI\\ into stars is the ratio of the $x$ axis values to the $y$-axis values in Figure \\ref{fig-nhisfr}, or approximately $10^{11}$ years (see also Bigiel et al. 2010).  The \\HI\\ distributions are clumpy with giant clouds that appear associated with star formation. These clouds are most obvious in the inner regions where the star formation rates are highest, although faint \\HI\\ clouds can also be seen in the outer parts where past-average star formation rates are very low.  Observable clouds typically have masses of around $10^7\\;M_\\odot$ or larger. A few faint clouds in the outer part of DDO 86 with masses of several $\\times10^6\\;M_\\odot$ were discussed as possible future sites of star formation.  In general there are no spiral arms or systematic streaming motions in \\HI. The velocity fields reflect normal rotation curves except for NGC 4163 and DDO 53.  The most extended \\HI\\ rotation curves are available for DDO 86, DDO 133, and IZw 115. The first two have rising inner rotation curves that flatten out at the break radius in the $V$-band profile. The third has a rising rotation curve in the inner part and a falling rotation curve in the outer part, with no break in the exponential profile for comparison. There could be a small companion galaxy near IZw 115 that affects the rotation curve in the outer part. The coincidence in two cases between the turnover point in the \\HI\\ rotation curve and the break point in the photometric radial profile is interesting but maybe not significant. In our larger sample (Hunter \\& Elmegreen 2006), there was no correlation between the photometric break radius and the rotation curve turnover radius: only one case had these features at the same radius, 5 had the break at a larger radius than the turnover and 4 had the break at a smaller radius than the turnover. In larger spiral galaxies, the exponential break radii are typically at $\\sim2-5R_{\\rm D}$ (Pohlen \\& Trujillo 2006) and the turnover radii are at $\\sim1R_{\\rm D}$, so there is no coincidence of radii for spirals either.  The \\HI\\ rotation speeds in these dwarf galaxies are typically low: in the three cases with extended rotation curves, they peak at 40 km s$^{-1}$ for DDO 86 and DDO 133, and at 15 km s$^{-1}$ for IZw 115. The \\HI\\ velocity dispersions are approximately constant with radius at around 10 km s$^{-1}$.  For IZw 115, the average rotation speed is comparable to the velocity dispersion inside $3R_{\\rm D}$ and possibly in the outer regions too where the rotation curve drops.  At the peak of the rotation curve in IZw 115, the ratio of the rotation speed to dispersion is remarkably low, $\\sim1.4$.  Such slow rotation gives IZw 115 very little dynamical integrity, yet it has a near-perfect exponential disk for $8R_{\\rm D}$.  Slowly rotating galaxies should be easily perturbed by interactions with other galaxies and dark matter concentrations, by clumping instabilities in their disks, and by gas blowout during star bursts. This sensitivity to asymmetric forcing and rapidly-changing potentials makes the long-range exponential structure of the disks difficult to understand, unless it results from such forcings.  The uniform rotation observed in 2 out of 5 cases suggests that their recent lives have been quiescent, at least for the last several rotation periods in their outer regions. The rotation period at the outer radius of detection for the \\HI\\ rotation curve in DDO 133 is 0.9 Gyr. For DDO 86, it is 1.9 Gyr. %and for IZw 115 it is 2.0 Gyr. The SFR ratios shown in Figure \\ref{fig-sfrrat} are roughly consistent with this long term quiescence.  The maps of velocity dispersion, ``mom2'', in Figures \\ref{fig-d53vf} to \\ref{fig-izw115vf} show an inverse correlation with the maps of column density, ``mom0,'' in the sense that higher column densities correspond to lower velocity dispersions.  We found this before for NGC 2366 (Hunter et al. 2001).  Figure \\ref{fig-m0m2} shows this anti-correlation better. Each panel shows a map of mom0 in red and mom2 in green. Overlapping regions turn out yellow. The maps show mostly red and green regions with only a little yellow. This means that the dispersion is high where the column density is low, and vice versa. An anticorrelation is roughly consistent with pressure uniformity in the sense that low density regions have higher turbulent speeds. Piontek \\& Ostriker (2005) found this in simulations of the magneto-rotational instability.  Trujillo \\& Bakos (2010) looked for a correlation between possible inflection points in the \\HI\\ velocity dispersion and exponential break radii in the disks of local spiral galaxies. They were considering that the breaks could result from gas disk flaring, but found nothing. We see nothing like that here either.  \\subsection{Gravitational Stability}  The \\HI\\ appears to be gravitationally stable by itself. The ratios of the gas surface densities to the Toomre (1964) threshold densities are 0.6 or less in DDO 86 and DDO 133, 0.1 or less in NGC 4163, and 0.3 in IZw 115. There are stars mixed with the gas, however, and the stellar velocity dispersions are probably not much larger than the gaseous dispersions, considering the low rotation speeds. This means that stars contribute to the instability in approximate proportion to their mass (Jog \\& Solomon 1984).  We have estimated the mass surface densities of stars in our galaxies using the $V$-band radial profiles and the $B-V$ colors, where available. The calibration comes from $M/L$ tables in Bell \\& de Jong (2000).  We find that the stellar mass surface densities in the outer parts of our galaxies are comparable to the gas surface densities, or slightly smaller, to within a factor of a few. This can also be derived from the star formation rates given in Table \\ref{tab-sfrs}, which are $\\sim10^{-4.8}\\;M_\\odot$ yr$^{-1}$ kpc$^{-2}$ in the $V$-band at $\\mu_{\\rm V}=29.5$ mag arcsec$^{-2}$. Multiplying this by $10^{10}$ years and converting to pc$^{-2}$, the result is a total stellar surface density today of $\\sim0.2\\;M_\\odot$ pc$^{-2}$ in the far outer parts. Figures \\ref{fig-d53sb} to \\ref{fig-izw115sb} give $\\Sigma_{\\rm gas}$ there, and the average is about $1\\;M_\\odot$.  Thus gravitational forcing from stars may be $\\sim20$\\% of that from gas during the growth of an instability, leading to an upward correction to the surface density ratios by $\\sim20$\\%.  The correction would be larger in the inner regions, where the stellar contribution is larger. These corrections make DDO 86 and DDO 133 only slightly stable, which is typical for galaxies. For example,  Leroy et al. (2008) showed that the gas+star stability parameter equals about 1.5 for a wide range of radii inside dwarfs and spirals. The corresponding ratio of column density to critical column density is the inverse of this, 0.6. This is not much different than it is for DDO 86 and DDO 133.  NGC 4163 is not a good case for this stability analysis because the rotation curve is peculiar (\\S 4.3). IZw 115 could also have a peculiar rotation curve because of the companion galaxy, but even in the inner parts where the rotation curve is rising, the \\HI\\ column density is less than the critical value by a factor of 5. These inner parts are probably post-starburst, however, so they could well be very stable.  The situation with instabilities is more complicated than this because all of these galaxies have relatively thick disks, as inferred from the high ratios of HI velocity dispersion to rotation speed. Thickness $H$ decreases the in-plane accelerations from self-gravity at wavenumber $k$ by a factor of $\\sim(1+kH)^{-1}$ (Vandervoort 1970), which should be $\\sim0.5-0.7$ in our case.  This lowers the effective stability ratio $\\Sigma/\\Sigma^\\prime_{\\rm crit}$ by the same factor, since $\\Sigma^\\prime_{\\rm crit}=\\kappa\\sigma(1+kH)/\\pi G$ for $Q^\\prime=1$ with thickness corrections.  Here, $\\sigma$ is the velocity dispersion, $\\Sigma$ is the mass column density, and $Q^\\prime=\\kappa\\sigma(1+kH) / \\pi G \\Sigma$ is the thickness-corrected Toomre $Q$ parameter for epicyclic frequency $\\kappa$ (Elmegreen 2011). The ratios derived above for gas+stars in DDO 86 and DDO 133 without thickness corrections are comparable to those given by Leroy et al. (2008) for dwarf and spiral galaxies without thickness corrections. Therefore both are likely to have comparable stability ratios when thickness is included, which means both are slightly more stable than without the thickness correction. The ratios for the outer parts of dwarfs would still be smaller than they are for spirals.  Gravitational instabilities in dwarf irregular galaxies should look different than they do in large spiral galaxies. The characteristic length for spirals or clumps that form in a gravitational instability lies in the range $\\sigma^2/\\left(\\pi G \\Sigma^\\prime\\right)\\left(1\\pm \\left[1-Q^2\\right] \\right)^{-1}$. This is centered at $\\sigma^2/\\left(\\pi G \\Sigma^\\prime\\right)$, which should be compared with the disk scale length. In dwarfs with comparable velocity dispersions for the gas and stars, we can use the total mass column density for $\\Sigma^\\prime$ and include both components this way. In the outer parts of dwarf galaxies, $\\Sigma\\sim10\\;M_\\odot$ pc$^{-2}$, $\\Sigma^\\prime\\sim0.5\\Sigma$, and $\\sigma\\sim10$ km s$^{-1}$, so $\\sigma^2/\\pi G \\Sigma^\\prime =1.4$ kpc. For spiral galaxies, the length scale is about the same because the velocity dispersion for the dominant stars is higher by a factor of $\\sim3$, and the mass column density is higher by a factor of $\\sim10$. Dwarf galaxies have smaller exponential scale lengths than spirals, however. Three of the dwarfs studied here have $R_{\\rm D}<0.7$ kpc and two others have $R_{\\rm D}\\sim2$ kpc (Table \\ref{tab-disk}). In spiral galaxies, $R_{\\rm D}$ ranges between $\\sim1$ and 5 kpc (van der Kruit 1987), which is larger than in dwarfs by a factor of 2 or more.  This difference means that the characteristic length of a significantly gravitating structure in a large spiral galaxy readily fits into the disk, allowing coherent structures like spiral waves to form. Most dwarfs are too small for this.  In addition, the bright parts of dwarf galaxy disks occur in the rising parts of their rotation curves, but the main parts of spiral galaxy disks are in the flat parts of their rotation curves. Gravitational perturbations therefore tend to grow as unsheared clumps in the dwarfs, but they produce spiral wakes and long sheared arms in spirals.  Some dwarf galaxies have spirals, but this situation is rare (Schombert et al. 1995). A dwarf elliptical in the Virgo cluster, IC 3328, has a weak spiral, but it may be confined to a thin sub-component of the main stellar disk where the gravitational length is short (Jerjen et al. 2000; De Rijcke et al. 2003; Lisker \\& Fuchs 2009). Other dwarf spirals in Virgo were observed by Barazza et al. (2002). %!! added ref above  The same considerations apply to the far-outer parts of spiral galaxies, where $\\Sigma$ is comparable to that in dwarfs. Then the characteristic length for instabilities may be 10 times higher than it is in the inner spiral disk, and locally much larger than the scale length. Thus the outer parts of spiral disks should have difficulty generating self-gravitating waves on their own, like dwarf galaxies. However, the outer disks of spiral galaxies can have spiral waves that propagate there from the inner disk (Bertin \\& Amorisco 2010) or were resonantly excited by perturbations from the inner disk (Debattista et al. 2006). Dwarf galaxies have neither source of waves and should produce only clumps of self-gravitating gas.  Disk stability is also relevant to the origin of double exponential profiles, according to the model by Debattista et al. (2006). This model involves angular momentum transfer around bars and spirals. Debattista et al. (2006) showed that the inner parts of barred galaxies with single exponential profiles get flatter with time, out to the outer Lindblad resonance of the spiral that is driven by the bar. The original exponential disk then shows up as the outer steep part of the double exponential, and the bar-modified part is the inner flat region. The degree of flattening depended on $Q$ in these models. Disks that were more gravitational stable, with higher $Q$, did not evolve as much as less stable disks. This result suggests that dwarf galaxies that are weakly self-gravitating may not be responsive to bar torques that modify a single exponential into a double exponential.  Our galaxies with the highest ratios of $\\Sigma_{\\rm gas}/\\Sigma_{\\rm crit}$, namely DDO 86 and DDO 133, also have the most obvious breaks in their exponential profiles. DDO 133 is a barred galaxy too. However, the break radius in DDO 133 is at the end of the bar and not the end of a spiral as in the Debattista et al. model. Also in DDO 86, the star formation region and inner bright core in Figure \\ref{fig-d86} is elongated in the east-west direction, reminiscent of a bar.  The end of this elongation corresponds to the break radius. Similarly in DDO 53, which is the only other case here with a double exponential, there is an elongation of the main disk star-forming region along a position angle of $\\sim150^\\circ$, which is transverse to the major axis. The deprojected length of this elongation is comparable to the break radius. There is not enough \\HI\\ data to determine $\\Sigma_{\\rm gas}/\\Sigma_{\\rm crit}$ for this galaxy.  These observations suggest that the inner exponentials of our galaxies were made flatter than the outer exponentials by mass transfer connected with a bar. There are no spirals here, so the bar may be doing here what the spirals did in Debattista et al. (2006).  We looked for a correlation between photometric breaks and the presence of bars in our larger sample (Hunter \\& Elmegreen 2006).  Out of 22 dwarf galaxies with double exponentials of the type discussed here, 8 have some evidence for a bar (even if they are not classified as Hubble bar types), and the rest do not. DDO 133, studied here, was one of the barred cases with a double exponential recognized in that previous study. Among the barred dwarfs, the ratio of the break radius to the end of bar radius is $1.8\\pm0.6$ on average. Since there are breaks without bars and bars without breaks, we do not see a clear correlation between the two at this time.  \\subsection{Two Peculiarities}  We summarize this discussion by highlighting two peculiar things about the present observations: the extraordinarily regular exponential disks in dwarf irregular galaxies, and the continuity of star formation throughout these disks and over time. Exponential disks are made in part by cosmological infall (Freeman 1970; Fall \\& Efstathiou 1980; van der Kruit 1987) and subsequent dynamical torques from self-gravitating disk structures. There is no regularity to these structures in dwarfs and no systematic presence of bars or spirals. There are only clumps of star formation and maybe some old disk stars mixed in with those clumps. It is conceivable that the clumps of star formation make or maintain the exponential disk, as in numerical simulations of equally clumpy galaxies in the early universe (Bournaud et al. 2007). Or, there could be bar-like dark matter halos (Bekki \\& Freeman 2002; Tutukov \\& Fedorova 2006) that redistribute angular momentum, while also driving in gas to fuel occasional low-level bursts (Hunter \\& Elmegreen 2004).  As for star formation itself, this complex process involving mostly gas seems to follow the underlying exponential disk of old stars better than one would expect from such a stochastic process. Leroy et al. (2008) found the same thing for dwarfs and the outer parts of spiral galaxies. Molecular cloud distributions follow the stellar exponential even when the molecules are an extremely low fraction of the gas mass (Schruba et al.\\ 2011). How do the molecules know to do this? In the current paradigm, star formation follows only the molecules, and molecular abundances depend mostly on the pressure (e.g., Bigiel et al. 2008), but why does this end up reinforcing the underlying exponential disk? There is a purely radial dependence to the cloud and star formation rates, in addition to a gas dependence, as in the dynamical law of star formation, where the rate is determined by the orbit time (Kennicutt 1998). Bigiel et al. (2010, Figure 7) also found that for a given \\HI\\ column density, the star formation rate decreases with increasing radius in spiral and dwarf galaxies.  The problem with this is that orbital timescales usually vary with radius as a power law, not an exponential.  Ostriker et al. (2010) proposed that the star formation rate scales with the total gas surface density multiplied by the square root of the midplane density from stars and dark matter. They get good agreement between this prediction and the radial profiles of star formation in several spiral galaxies.  The present results suggest that if this quantity applies to dwarf irregulars, then it should also be proportional to the column density of existing stars so that the color profile does not change much over at least a Gyr.  In fact they find the star formation rate per unit gas mass proportional to the surface density of existing stars and dark matter in the case where the vertical velocity dispersion is constant (because then the midplane density scales with the square of the stellar and dark matter column density).  This matches our result fairly well because the gas column density varies weakly with radius.  Another model is that in highly stable regions, molecular cloud formation and star formation is somehow triggered by field stars. Perhaps Type Ia supernovae are important, or the turbulence that is driven by these supernovae and other field-star pressures. Star formation would then follow the underlying stellar density as long as there is enough gas to compress. Since the rate is approximately constant over time, sequential triggering by young star pressures could also be important.  Then star formation would move around as illustrated by Gerola et al. (1980).  A way to check these models would be to find triggering sources near young stars.  Also, the star formation rate should be more tightly correlated with the stellar column density than with the gas column density, as long as there is enough gas to be triggered.  The break radii in exponential disks could be a clue to these processes if the properties of the gas and cosmic rays near regions of star formation were found to differ inside and outside of this radius. In the present sample of galaxies, we found no clear correlation between the existence of a break and other properties, such as bars, spirals, color jumps, or the turnover in the rotation curve. We found no such correlations in our bigger sample either (Hunter \\& Elmegreen 2006)."
    },
    "1107/1107.0611_arXiv.txt": {
        "abstract": "We investigate the relation between the star formation history and the evolution of the double-degenerate (DD) population in the thin disc of the Galaxy, which we assume to have formed 10 Gyr before the present. We introduce the use of star-formation contribution functions as a device for evaluating the birth rates, total number and merger rates of DDs. These contribution functions help to demonstrate the relation between star-formation history and the current DD population and, in particular, show how the numbers of different types of DD are sensitive to different epochs of star formation.  Analysis of the contribution functions given by a quasi-exponentially decline in the  star-formation rate shows that star formation from  0 to 8 Gyr after thin-disc formation dominates the present rates and total numbers of He+He DDs and CO+He DDs. Similarly, the current numbers of CO+CO and ONeMg+X DDs come mainly from early star formation ($<6$ Gyr) , although star formation from 4 to 8 Gyr continued to contribute more CO+CO than He+He DDs. The present birth-rates for CO+CO and ONeMg+X DDs are strongly governed by recent star formation ({\\it i.e.} $8 - 9.95$ Gyr). Star formation from $<7.5$ Gyr does {\\it not} contribute to the present birth rates of CO+CO and ONeMg+X DDs, but it has a distinct contribution to their merger rates.  We have compared the impact of different star-formation models on the rates and numbers of DDs and on the rates of type Ia (SNIa) and core-collapse supernovae (ccSN). In addition to a quasi-exponential decline model, we considered an instantaneous (or initial starburst) model, a constant-rate model, and an enhanced-rate model.  All were normalised to produce the present observed star density in the local thin disc. The evolution of the rates and numbers of both DDs and SNIa are different in all four models, but are most markedly different in the instantaneous star-formation model, which produces a much higher rate than the other three models in the past, primarily as a consequence of the normalisation.  Predictions of the current SNIa rate range from  $\\approx2$ to $5\\times10^{-4}$ yr$^{-1}$ in the four models, and are slightly below the observed rate because we only consider the DD merger channel. The predicted ccSN rate ranges from 1.5 to 3 century$^{-1}$, and is consistent with observations. ",
        "introduction": "\\label{sec_introduction}  We know little about the star formation (SF) history of the Milky Way disc. Of the few ways available to explore the SF history, studying the stellar age distribution by the observation of a sample of different types of stars and comparing observed and synthetic colour-magnitude diagrams would be the most favored approach. The simplest view of the SF history would be a single star burst taking place at the formation of the disc. Substantial evidence, however, shows the disc has experienced a more complicated history, including repeated star bursts, an exponential declining SF or extended periods of enhanced SF \\citep{Majewski93,Rocha-Pinto00} or some combination of all of these.  The present SF rate is better understood. \\citet{Smith78} concludes the SF rate in the Galaxy is perhaps 5 $M_{\\odot}$yr$^{-1}$, of which 74\\% take place in spiral arms, 13\\% in the interarm region, and 13\\% in the galactic center. Observations of Lyman continuum photons from O stars in giant H II regions give the SF rate in the disc to be 4.35 $M_{\\odot}$yr$^{-1}$. A SF rate $\\approx$3.6 $M_{\\odot}$yr$^{-1}$ was suggested by \\citet{McKee89} from an analysis of thermal radio emission in HII regions around massive stars. The measurement of the mass of $^{26}$Al in the Galaxy implies a SF rate $\\approx$4 $M_{\\odot}$yr$^{-1}$ \\citep{Diehl06}. The derivation of the SF rate is strongly associated with the initial mass function (namely the distribution of the mass of the new-born stars) and the birth rate of core collapse supernovae (type Ib/c and type II).  Double degenerates (DDs) are a type of exotic binary consisting of two white dwarfs. As the evolved remnants of stars, the formation of close DDs requires their progenitors to have undergone a mass transfer stage, either common envelope or stable Roche lobe overflow phase. Ultra-compact DDs are especially interesting since not only can they be observed as active sources of electromagnetic radiation, {\\it e.g.} AM Canum Venaticorum stars, \\citep{Cropper98,Ramsay05,Nelemans04}, but also the final products of DD evolution may include type Ia supernovae \\citep{Yungelson00}). Theoretical results \\citep{Evans87,Nelemans01b,Ruiter10,Yu10} suggest that close DDs should also be significant sources of gravitational wave radiation, and be detectable by the space gravitational wave (GW) detector LISA (Laser Interferometer Space Antennae).  The census of DDs therefore is important because they allow us to test the endpoints of stellar evolution. We need to know the present-day numbers, birth rates, merger rates and galactic distribution of DDs, and also the distribution of their properties ({\\it i.e.} mass, orbital separation, chemical abundance and age). To interpret these, we need to investigate how the history of star formation in the Galaxy affects the DD population. We can achieve this using a population synthesis approach.  In this paper, we study the correlation of the SF history and the evolution of DD population. It may help interpret the distribution of the physical properties of DDs as a function of the SF history. In \\S \\ref{sec_population}, we describe the approach to simulate a population of the DDs in the thin disc and discuss the details of how a quasi-exponential declining SF rate influences the birth rate, merger rate and total number of DDs. In \\S \\ref{sec_sfimpact}, we show the distributions of some important physical parameters of the DDs from different epochs of SF. In \\S \\ref{sec_comparison}, we compare the impact of different SF models on the rates and numbers of DDs. We discuss the results for supernovae rates in different SF models in \\S \\ref{sec_snia}. Some points are discussed in \\S \\ref{sec_summary} and conclusions are drawn in \\S \\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion}  In this paper we have investigated the relation between the evolution of birth rates, merger rates and numbers of different types of DD and the SF history in the thin disc of the Milky Way Galaxy. By analysing how a quasi-exponentially declining SF rate influences the rates and observed numbers of DDs with respect to the evolution of the disc and SF, we find that SF between 0 and $\\sim$8 Gyr dominates the present rates and total numbers of He+He DDs and CO+He DDs. Similarly, the current numbers of CO+CO and ONeMg+X DDs mainly come from early SF, although the later SF ({\\it e.g.} 4 to 8 Gy) contributes more CO+CO than it does for He+He which are the two largest DD population in the thin disc.  However, the present birth and merger rates of CO+CO and ONeMg+X DDs are strongly governed by the recent star formation ({\\it i.e.} 8 to 9.95 Gyr). SF before $\\approx$7.5 Gyr does {\\it not} contribute to the present birth rates of CO+CO ang ONeMg+X DDs, but it has some contribution to the merger rates. More importantly, with the SF history related population synthesis approach in this paper, not only are we able to determine the rates and numbers of DDs, but we can also obtain the distributions of properties of current DD population from different stages of SF.  We have compared the impact of different SF models, namely the instantaneous, the constant, the enhanced, and the quasi-exponential SF, on the rates and numbers of DDs and the rates of two types of supernovae. The evolution tracks of the rates and numbers from the four models are quite different. A distinct difference can be found between the instantaneous model and the other three continuous SF models. This model gives historical rates and numbers obviously higher than the other three models. However, the present rates of DDs from this model are apparently lower than the other three models.  In addition to the DDs, we have calculated the rates of SNIa and ccSN. The evolution of the rates of SNIa is basically similar in the four SF models, but the instantaneous model can produce a higher rate in the past because of the very high SF rate at the formation of the thin disc. The present rates of SNIa are 2.19, 4.92, 5.07, and 4.34 $\\times10^{-4}$ yr$^{-1}$ for the instantaneous, the constant, the enhanced, and the quasi-exponential SF models respectively. The rates of ccSN from all four SF models, $\\approx$1.5 to 3 century$^{-1}$, are consistent with the observations."
    },
    "1107/1107.4612_arXiv.txt": {
        "abstract": "The upcoming generation of cosmic microwave background (CMB) experiments face a major challenge in detecting the weak cosmic $B$-mode signature predicted as a product of primordial gravitational waves. To achieve the required sensitivity these experiments must have impressive control of systematic effects and detailed understanding of the foreground emission that will influence the signal. In this paper, we present templates of the intensity and polarisation of emission from one of the main Galactic foregrounds, interstellar dust. These are produced using a model which includes a 3D description of the Galactic magnetic field, examining both large and small scales. We also include in the model the details of the dust density, grain alignment and the intrinsic polarisation of the emission from an individual grain. We present here Stokes parameter template maps at $150$ GHz and provide an on-line repository\\footnote{\\url{http://www.imperial.ac.uk/people/c.contaldi/fgpol}} for these and additional maps at frequencies that will be targeted by upcoming experiments such as \\ebex, \\spider\\ and \\spt pol. ",
        "introduction": "The next round of cosmic microwave background (CMB) experiments are all targeting measurements of CMB polarisation.  The CMB polarisation field can be decomposed into curl-free $E$-modes and curl-like $B$-modes (see for example \\cite{1997PhRvD..55.7368K}), however to date only $E$-modes have been observed \\citep{DASI,CBIpol,2007ApJ...660..976S, CAPMAP,2008ApJ...684..771B, DASI3yr,BoomerangEE, 2006ApJ...647..833P,2007ApJ...665...55W, 2007ApJS..170..335P,2009ApJ...705..978B, 2010ApJ...711.1123C, 2010arXiv1012.3191Q}. Current work centres on the search for the tiny amplitude $B$-modes since a key prediction of inflation is the generation of a stochastic background of gravitational waves, which on large angular scales are the only contribution to the CMB $B$-mode component. Experiments involved in this search include \\ebex\\ \\citep{2010SPIE.7741E..37R}, \\spider\\ \\citep{2010SPIE.7741E..46F}, \\spt pol~\\citep{2009AIPC.1185..511M}, \\piper\\ \\citep{2011AAS...21823301L}, \\abs\\ \\citep{2010arXiv1008.3915E}, \\act pol \\citep{2010SPIE.7741E..51N}, \\polar\\ \\citep{1998ApJ...495..580K}, \\polarbear\\ \\citep{2010arXiv1011.0763T} and \\brain\\ \\citep{2008arXiv0805.4527C}. We focus on providing templates at frequencies being targeted by \\ebex , \\spider\\ and \\spt pol. Galactic foreground emission is expected to contribute significantly to the polarised microwave emission across the sky, and may well dominate the CMB gravitational-wave signal at all frequencies.  In order to achieve their scientific goals, forthcoming CMB polarisation experiments require in-depth knowledge of this polarised Galactic foreground emission. The \\planck\\ mission \\citep{2006astro.ph..4069T} will provide maps of the polarisation of interstellar dust, allowing tests of the structure of the Galactic magnetic field. It should also provide an insight into grain alignment mechanisms.  Of utmost importance will be the accurate separation of foreground emission from the CMB signal. Component separation has been considered in, for example, \\citet{1994ApJ...424....1B, 2006ApJ...641..665E, 2007ApJ...665..355K, 2009MNRAS.392..216S}. In order to test and assess methods for separating out the contribution realistic foreground templates will be required. Since however, few data exist at this time at the observing frequencies in which the upcoming experiments are operating, one must resort to modeling of foreground emission by extrapolating the information from existing data.  Furthermore, experiments which observe portions of the sky close to the Galactic plane will find the foreground emission is very bright in comparison to the CMB. The presence of this bright emission in the data may affect the performance of the observation and the analysis strategy an experiment uses. Another important role of foreground modeling, therefore, is informing the planning and proposal stage of any experiment.  Unfortunately, these foregrounds are poorly constrained by current data and poorly understood, particularly above around $90$\\,GHz, where the CMB emission is strongest and where many new CMB experiments will operate.  At these frequencies, the foreground emission is expected to be dominated by thermal emission from interstellar dust. A review of the basic physical processes whereby aligned dust grains generate polarisation is given in \\citet{2003rrda.book..255L}. Such emission is known to be polarised both through direct measurements \\citep{2005A&A...444..327P, 2007ApJ...665..355K, 2004A&A...424..571B,Bierman:2011hq} and through observations of the polarisation of starlight \\citep{1996ApJ...462..316H, 2002ApJ...564..762F}.   This polarisation arises due to the presence of a magnetic field in the Galaxy. The dust grains are generally non-spherical, and preferentially emit radiation polarised along their longest axis. Mechanisms exist which align these grains with this axis perpendicular to the Galactic magnetic field, leading to net linear polarisation.    \\begin{figure*} \\begin{center}\\begin{tabular}{c} \\makebox[3in][c]{\\includegraphics[height=7cm, width=7cm,angle=0,trim=0cm 0cm 0cm 0cm,clip]{galaxy_top}\\includegraphics[height=7cm, width=7cm,angle=0,trim=0cm 0cm 0cm 0cm,clip]{galaxy_lsa_top}} \\\\ \\end{tabular} \\caption[]{Cartesian projection of the BSS (left panel) and LSA (right panel) large-scale magnetic field models, showing a slice through the Galactic plane observed along the positive $z$-axis. The filled circle represents the position of the Sun in relation to the Galactic centre. The alignment and magnitude of the magnetic field are shown as headed ticks with the $B_0=2$ $\\mu$G scale represented at the top of each panel. The red solid lines are the density contours, in steps of 0.1, of the dust density model $n_d(r,z)$ (see Section~\\ref{sec:dustdens}) used in the line-of-sight integration. The dust density is normalised to 1 at the Galactic centre.} \\label{fig:lsmodel} \\end{center} \\end{figure*}  Polarised foregrounds also include polarised emission from synchrotron. Synchrotron emission, generated by the gyration of cosmic ray electrons in the Galactic magnetic field, is intrinsically polarised and constitutes the main polarised foreground at lower frequencies \\citep{2007ApJS..170..335P}.  However, emission from thermal dust dominates synchrotron emission at the frequencies considered in this paper and therefore we concentrate on modeling emission from only thermal dust. Foreground radiation also includes free-free and spinning dust emission, however we assume both of these signals to be unpolarised and so do not consider them further here. Evidence for this has emerged recently, with \\citet{2011arXiv1108.0205M} showing that free-free emission is unpolarised, setting an upper limit on the free-free polarisation fraction of 3.4\\% at the $2\\sigma$ level. They also show that spinning dust emission has a low polarisation fraction. Upper limits on the polarisation fraction of spinning dust emission in molecular clouds have been obtained by \\citet{2011arXiv1108.0308D} and \\citet{2011ApJ...729...25L}, who find low levels of polarisation. If there is a similar level at higher Galactic latitude, then this foreground is unimportant in terms of component separation.  The model in this paper was introduced by \\cite{Danielthesis:2009} and first applied in~\\citet{2011arXiv1102.0559O} for the purpose of studying the impact of polarised foregrounds on Spider's ability to detect $B$-mode polarisation.  Here we give a more detailed explanation of the model and present a number of full-sky template maps at various frequencies.  \\archeops\\ \\citep{2004A&A...424..571B} and {\\sc BICEP} \\citep{Bierman:2011hq} have made the highest signal--to--noise maps of the dust polarisation at low Galactic latitudes and have examined the properties of the polarisation fraction and angle. These results cannot be relied upon to calibrate large scale models of the polarisation at higher latitudes since the polarisation properties near the Galactic plane will depend on complex structure that is not included in models such as the one presented in this work. Here, both polarisation amplitude and angle are modeled internally and our templates are scaled such that the polarisation fraction corresponds to a nominal value when averaged over the maps with the Galaxy masked out. The \\archeops\\ and {\\sc BICEP} maps have not been made public and a quantitative comparison of our templates with these observations is not possible. \\citet{AA..526..A145..2011} have developed a similar model of both the polarisation of thermal dust and synchrotron radiation and compared with the \\WMAP\\ K-band and \\archeops\\  353GHz data, however they have not released any templates based on their model.   This paper is organized as follows.  We begin in Section~\\ref{sec:total} by describing the model of the dust total intensity.  We focus on modeling the Galactic magnetic field (which is made up of a large-scale field from the Galactic disk and a small-scale field due to turbulence in the interstellar plasma) in Section~\\ref{sec:galacticfield}.  We include in Section~\\ref{sec:dustdens} a description of the large-scale dust density, as well as details of the treatment of dust grain alignment and the intrinsic polarisation of the emission from an individual grain. In Section~\\ref{sec:stokes} we show how these fields are modeled in three-dimensions and are used to perform a line-of-sight integral to the centre of each pixel to form the final Stokes parameter maps.  In Section~\\ref{sec:scales} we discuss the effective resolution of the maps produced using our model and the relevant physical scales introduced by various effects. In Section~\\ref{sec:maps} we describe the final template maps and summarise our conclusions in Section~\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude} We have described a model of the polarised foreground that we expect to observe due to emission from dust within our Galaxy. The model uses a three-dimensional model of the Galactic magnetic field and dust field and integrates along the line-of-sight to each \\healpix\\ pixel to obtain a polarisation amplitude and angle. This information is combined with total intensity, FDS derived template maps at different frequencies to obtain a complete, polarisation template of foreground emission by interstellar dust.  We have concentrated on two popular models for the structure of the large-scale structure of the Galactic magnetic field, namely, the BSS and LSA models. The parameters for the BSS model have been calibrated directly from measurements of the strength of the Galactic magnetic field. In the LSA case we have employed the parameters obtained by \\cite{2007ApJS..170..335P} in fitting to the \\WMAP\\ observations. We calculate the polarisation alignment {\\sl internally} to our model in both cases since there is not sufficient external information on polarisation angles at resolutions relevant in this work. Some differences exist between the BSS and LSA derived templates but these are mostly at low Galactic latitudes away from the Galactic centre and as such experiments targeting small areas at high Galactic latitudes will not be sensitive to the differences. The differences do indicate however that a more accurate model of the Galactic magnetic field is required to produce realistic polarisation templates for low Galactic latitudes.  In the future, the \\planck\\ mission will provide an important test of Galactic magnetic field models through detailed characteristion of galactic foregrounds.  We have developed a one dimensional approximation of the stochastic, turbulent, small-scale component of the field for obtaining full-sky templates. A full three dimensional realisation of the turbulent component can be used to obtain higher resolution templates for smaller patches of the sky.  In future work we will be extending the model to include synchrotron emission to form a complete picture of foreground emission relevant for polarisation experiments. Other developments will be required to increase the fidelity of the templates on small scales. These include the addition of a stochastic, small-scale density field to model small-scale structure in the density. In full ``3D'' calculations this will require the generation of an additional three dimensional, turbulence realisation which is correlated to the small-scale magnetic field. In addition, it would be useful to develop a simple model for the correlation of both realisations with the small-scale structure in the FDS derived total intensity templates.  There is significant freedom in the parameters defining the small-scale structure in the templates. Experiments targeting small angular scales over small patches of the sky will be most sensitive to variations in the parameters and also to the stochasticity of the structure in the templates. Further Monte Carlo explorations of the variation in the maps is therefore warranted to quantify the impact of foregrounds on future sub-orbital experiments. As part of future work we will generate large ensembles of random realisations of the templates on small patches of the sky for the purpose of Monte Carlo studies.  The maps obtained from this model are available for use and can be downloaded from an on-line repository\\footnote{\\url{http://www.imperial.ac.uk/people/c.contaldi/fgpol}}."
    },
    "1107/1107.0946_arXiv.txt": {
        "abstract": "We study the excitation and damping of tides in close binary systems, accounting for the leading order nonlinear corrections to linear tidal theory.  These nonlinear corrections include two distinct physical effects: three-mode nonlinear interactions, i.e., the redistribution of energy among stellar modes of oscillation, and nonlinear {\\it excitation} of stellar normal modes by the time-varying gravitational potential of the companion.  This paper, the first in a series, presents the formalism for studying nonlinear tides and studies the nonlinear stability of the linear tidal flow. Although the formalism we present is applicable to binaries containing stars, planets, and/or compact objects, we focus on non-rotating solar type stars with stellar or planetary companions. Our primary results include the following: (1) The linear tidal solution almost universally used in studies of binary evolution is unstable over much of the parameter space in which it is employed.  More specifically, resonantly excited internal gravity waves in solar-type stars are nonlinearly unstable to parametric resonance for companion masses $M'\\ga 10-100 M_\\Earth$ at orbital periods $P\\approx 1-10\\trm{ days}$.  The nearly static ``equilibrium'' tidal distortion is, however, stable to parametric resonance except for solar binaries with $P\\la 2-5 \\trm{ days}$. (2) For companion masses larger than a few Jupiter masses, the dynamical tide causes short length scale waves to grow so rapidly that they must be treated as traveling waves, rather than standing waves.  (3) We show that the global three-wave treatment of parametric instability typically used in the astrophysics literature does not yield the fastest growing daughter modes or instability threshold in many cases. We find a form of parametric instability in which a single parent wave excites a very large number of daughter waves ($N\\approx 10^3[P/\\trm{ 10 days}]$ for a solar-type star) and drives them as a single coherent unit with growth rates that are a factor of $\\approx N$ faster than the standard three-wave parametric instability. These are local instabilities viewed through the lens of global analysis; the coherent global growth rate follows local rates in the regions where the shear is strongest. In solar-type stars, the dynamical tide is unstable to this collective version of the parametric instability for even sub-Jupiter companion masses with $P \\la$ a month. (4) Independent of the parametric instability, the dynamical and equilibrium tides excite a wide range of stellar p-modes and g-modes by nonlinear inhomogeneous forcing; this coupling appears particularly efficient at draining energy out of the dynamical tide and may be more important than either wave breaking or parametric resonance at determining the nonlinear dissipation of the dynamical tide. ",
        "introduction": "\\label{sec:intro}  In binary systems at small separation, the differential gravity from one body can induce a significant tide in the other. As long as the tidal potential is time-varying, the tide and orbit exchange energy and angular momentum. Dissipation damps away the variable component of the tide, leading to an evolution of both the orbit and the spins of the bodies. The final state of such dissipation, if it is allowed to continue, is a circular orbit, with the spin and orbital angular momentum aligned, and the bodies synchronized with the same side always facing each other.  The rate at which the orbit and spins evolve depends on two quantities: the amplitude of the tide and how rapidly the energy in the tide is dissipated (``tidal friction\"). By analogy with a damped, driven oscillator, energy dissipation at a rate $\\dot{E}$ acting on a tide with frequency $\\omega$ and energy $E$ induces a phase lag $\\delta \\sim \\dot{E}/(\\omega E)$ between the time of maximum tidal acceleration and high tide. This phase lag is often written in terms of the tidal quality factor $Q\\simeq 1/\\delta$, where large $Q$ implies little dissipation and vice versa. First modeled in the classic treatment of \\citeauthor{Darwin:1879} (1879; see \\citealt{Goldreich:66} for applications in the Solar System), tidal friction is fundamental to our understanding of the origin, evolution, and fate of a wide variety of binary systems, including planets and moons in the solar system, extrasolar planets, solar-type binaries, and compact object binaries.  What is the physical origin of the `friction' which circularizes the orbits and synchronizes the spins of {\\it completely fluid} bodies in binary systems?  Most investigations treat the fluid motion as a linear perturbation of the background state, and focus on linear damping mechanisms such as radiative diffusion or viscosity from eddies in convection zones.  However, these applications of linear theory fail to explain many of the observed properties of binary systems, as we describe below.  This paper goes beyond the standard linear theory treatment in order to understand the effects of nonlinear interactions on tides in binary systems. Although aspects of this problem have been studied previously (e.g., \\citealt{Press:75,Kumar:96,Goodman:98,Barker:10}), we present a comprehensive formalism for understanding nonlinear tides in stars, planets, and even compact objects.  The ideas and formalism we present are general, but for concreteness all of our examples focus on tides in slowly rotating solar type stars with either stellar or planetary companions; in such stars internal gravity waves (g-modes) can be nearly resonant with the tide.  Our goal is ultimately to quantify the orbital and spin evolution of binaries accounting for nonlinear damping and excitation of tides.  In this first paper, however, we focus on determining the conditions under which standard linear theory is invalid because nonlinear instabilities drain energy out of the linear tide -- we leave calculations of the resulting Q-values and orbital evolution to future work.  Our study of tides is motivated by a number of observational puzzles, including:  1. {\\it Orbital circularization of solar binaries}---Binaries containing two solar-type stars, or a solar-type star and a gas giant exoplanet, are observed to circularize much more rapidly than expected from even the most detailed theoretical studies (e.g., \\citealt{Ogilvie:04, Meibom:05}). The observed circularization rates imply $Q\\sim10^5-10^6$, a factor of $100-1000$ times smaller than theoretical predictions. These observations test theories of tides in solar-type stars and gas giant planets over a range of orbital periods and amplitude of the tide.  2. {\\it Orbital decay of hot Jupiters}---The tidal Q inferred from observations of the circularization of solar-type binaries makes dire predictions for the orbital decay of hot Jupiters (Jupiter-size extrasolar planets with orbital periods $\\la 1$ week).  Tides raised in the star by the planet attempt to spin up the star, decreasing the orbital semi-major axis to conserve total angular momentum.  If the tidal Q for solar-type stars is really $10^5-10^6$, Darwin's theory predicts that all 5 Gyr old hot Jupiters inside 4 day orbits should have spiraled into their host (e.g., \\citealt{Ogilvie:07}); by contrast, there are $\\sim 100$ such planets observed.  The observed population of hot Jupiters could represent the tail end of a much larger parent population, most of which have decayed into their host stars \\citep{Jackson:09}.  Alternatively, the (near universal) practice of calibrating Q from one observation and applying it to another may be incorrect (e.g., \\citealt{Ogilvie:07}); it is certainly suspect without a deeper understanding of the underlying physics.  For example, such a calibration would be invalid if tidal excitation and dissipation depends on orbital frequency and/or the amplitude of the tide (i.e., the mass of the secondary).  In this paper, we show that this is indeed the case.  3. {\\it Radii of hot Jupiters}---Tidal heating has been invoked to explain the large observed radii of transiting hot Jupiters compared to evolutionary models \\citep{Gu:03, Arras:06, Arras:10}. Primordial eccentricity and asynchronous rotation will give rise to tidal heating, the duration of which is set by the value of $Q$. Steady state mechanisms, such as thermal tides \\citep{Arras:10}, may also be at work, for which the tidal heating rate is also dependent on $Q$. A question which has received less attention is the depth-dependence of tidal heating. In order for tidal heating to effectively bloat radii, the heat must be deposited in deep regions with long thermal times. Understanding the depth dependence requires detailed calculations of wave excitation and damping.  \\vspace{0.2cm}  The observational puzzles above, and the physical motivation that nonlinear effects can be important, led us to the present investigation. At first glance, it may be surprising that nonlinear effects are important, when basic estimates such as the height of the tide relative to the stellar radius are small. However, the importance of nonlinear effects is more subtle since wave amplitudes can become large in highly localized regions.  Consider, e.g., surface waves on the ocean which can travel great distances relatively undamped, but are susceptible to wave breaking as the depth decreases near shore.  In the next subsection we review the results from linear tidal theory in more detail; we also relate our nonlinear calculations to these standard linear theory results.  In \\S~\\ref{sec:structure} we summarize the structure of the remainder of this paper.  Throughout we presume a basic understanding of the physics of stellar g-modes and p-modes (see, e.g., \\citealt{Unno:89, Dalsgaard:03}).  \\subsection{Linear Theory Predictions for Solar-Type Stars}  The tidal response of a fluid is typically decomposed into ``equilibrium'' and ``dynamical'' components.  The simplest approximation for the tidal flow, the ``equilibrium tide,\" ignores fluid inertia, allowing the fluid to follow gravitational equipotentials. The equilibrium tide has large lengthscales, and hence is weakly damped by diffusive effects (such as turbulent viscosity due to convection), leading to circularization times much longer than are observed (e.g., \\citealt{Goodman:97}).  The ``dynamical tide\" refers to the resonant excitation of waves by the tidal acceleration.  This, together with linear mechanisms of dissipation, is the key physics in most {\\it linear} predictions of tidal dissipation.  In this paper, we calculate the nonlinear stability of both the equilibrium and dynamical tides.  The stability of the equilibrium tide in particular has rarely been considered in the literature (although see, e.g., \\citealt{Press:75}, \\citealt{Papaloizou:81}, and \\citealt{Kumar:98a} for some discussion along these lines).  It is, however, an {\\it a priori} promising source of enhanced orbital evolution given that the equilibrium tide contains most of the energy in linear theory.   Solar-type stars are composed of a stably stratified core and a neutrally stratified convective envelope.  Ignoring stellar rotation, tides in the convective envelope resemble the equilibrium tide. However, buoyancy in the core gives rise to short wavelength internal gravity waves that can have periods comparable to the orbital period of the binary; these waves have large amplitudes at the stellar center, but small amplitudes in the convection zone. Lastly, including both stellar rotation and fluid inertia, the core gravity waves are altered, and inertial waves restored by the Coriolis force are possible in both the radiative and convective zones.   The main linear dissipation mechanisms for solar-type stars are radiative diffusion in the core and turbulent viscosity due to convection in the envelope. \\citet{Terquem:98} and \\citet{Goodman:98} carried out detailed calculations for nonrotating stars. Both studies show that resonant gravity waves in the core are only weakly damped by radiative diffusion, giving $Q$ values far too large to be of interest. Goodman \\& Dickson sharpened the question. They showed that even if the train of gravity waves traveling inward from the core-envelope boundary -- where they are excited by the tidal potential -- were completely absorbed at the center the resulting dissipation would be $Q \\simeq 10^9\\ [P/(10\\ {\\rm days})]^{8/3}$ (based on their eq. [13]), too large by 3 orders of magnitude. Goodman \\& Dickson did not present detailed calculations of the damping of the gravity waves at the center, but argued that the waves would ``break,'' i.e., become highly nonlinear, based on an estimate of the gravity wave amplitude required for the waves to invert the stratification of the star.  This argument is supported by the numerical simulations of \\citet{Barker:10} and \\citet{Barker:11}.  Drawing on Goodman \\& Dickson's argument, \\citet{Ogilvie:07} suggested that the dissipation of gravity waves may be much less efficient for planetary companions because the amplitude of the gravity waves are much smaller when they are excited by a planet rather than a star (given the factor of $\\sim 1000$ mass difference between the two); in particular, they argued that the gravity waves would not overturn the stratification in the core and the efficient dissipation invoked by \\citet{Goodman:98} would not apply in the case of hot Jupiter systems (see also \\citealt{Barker:10}).  In this paper, we show that nonlinear damping of the dynamical tide can be important even when the stratification of the star is not overturned.  Linear theory thus does not apply even for solar-type stars with planetary companions.  Searching for a more efficient source of damping for tides in solar-type stars, \\citet{Witte:02} and \\citet{Ogilvie:07} realized that far shorter lengthscales, and hence enhanced damping, could be achieved by including the Coriolis force. Ogilvie \\& Lin showed that $Q \\sim 10^8$ can be achieved for solar-type stars based on the excitation and damping of short wavelength inertial waves in the convective envelope. This is, however, still not sufficient to explain the observed circularization of solar binaries.  \\subsection{Structure of this Paper} \\label{sec:structure}  The remainder of this paper is organized as follows.  In \\S~\\ref{sec:eom} we use perturbation theory to derive the nonlinear equations of motion and the key dimensionless coupling coefficients required to quantify the nonlinear couplings among stellar modes. Many of the details of these calculations are given in Appendix \\ref{sec:app:coef_in_amp_eqn}.  The results in \\S~\\ref{sec:eom} include two different ways of approaching the nonlinear problem.  In the first (\\S~\\ref{sec:method1}), we expand quantities relative to the unperturbed background state of the star.  In the second (\\S~\\ref{sec:method2} and Appendix \\ref{sec:app:alternative_eom}), we expand quantities relative to the linear tidal solution.  These two approaches are formally equivalent, but we will show that having both at our disposal is useful for understanding certain aspects of nonlinear tides in stars.  In Appendix \\ref{sec:properties_of_coefficient} we specialize to solar type stars and describe some of the key properties of the nonlinear coupling coefficients derived in \\S~\\ref{sec:eom} and Appendix \\ref{sec:app:coef_in_amp_eqn}.  Having derived the underlying equations of motion, we then review the results of linear theory (\\S~\\ref{sec:linear}) and summarize the key physical processes present in nonlinear tidal theory (\\S~\\ref{sec:nonlinear_tide}).  In \\S~\\ref{sec:stability_analysis} we present the nonlinear stability analysis that is at the heart of this paper: this allows us to determine the conditions under which the linear tidal flow is nonlinearly unstable to the parametric instability. We also find a `collective' parametric instability involving many modes that grows significantly faster than the standard three-mode parametric instability.  Section \\ref{sec:nonlinear} presents examples of the types of nonlinear instabilities that afflict stellar tides, including both resonant nonlinear coupling (parametric instability) and nonlinear inhomogeneous driving.   In \\S\\S~\\ref{sec:dyntide} and \\ref{sec:eqtide} we apply the results of \\S~\\ref{sec:stability_analysis} to solar type stars and assess the parametric stability of the dynamical and equilibrium tides, respectively.  In \\S~\\ref{sec:inhomog_driving} we study off-resonant non-linear excitation of both stellar g-modes and p-modes by the equilibrium and dynamical tides.  In \\S~\\ref{sec:global_vs_local} we assess when the nonlinear interactions are sufficiently strong to invalidate the assumption that tides excite global normal modes.  Finally, in \\S~\\ref{sec:conclusions} we summarize our results, discuss their implications for solar binaries and solar-type stars with hot Jupiter companions, and discuss key directions for future research.  The rigorous study of the synchronization and circularization of binary systems requires that one include the rotation of the star being tidally forced.  However, including even the simplest case of rigid rotation greatly complicates the analysis. While the centrifugal force may be ignored to a good approximation, inclusion of the Coriolis force changes the properties of waves with mode frequency smaller than rotation frequency, and new wave families appear. Specifically, eigenfunctions must be expanded in a sum over spherical harmonics, instead of a single harmonic for nonrotating stars. This greatly complicates the calculations and for simplicity we neglect Coriolis forces in this paper. ",
        "conclusions": "\\label{sec:conclusions}   This paper is an initial investigation into the importance of nonlinear fluid dynamics for the tidal evolution of close binary systems. We derive a formalism for including nonlinear interactions in tidal theory and describe the physical effects that arise from these nonlinearities.  While the ultimate goal is to understand how nonlinear effects may alter the rate of circularization and synchronization in binaries, such calculations are deferred to a future paper. In this paper, we instead focus on determining the conditions under which the standard linear theory approximation is (in)valid. We present detailed results for the tidal forcing of a sun-like star by a stellar or planetary companion, but the formalism we derive is more general and is applicable to nonlinear tides in stars, planets, or compact objects.  Most previous studies of tides have made the linear approximation, accounting for only linear order perturbations to the background state.  Our formalism includes the leading order nonlinear corrections. These corrections have two seemingly different physical origins: (1) \\emph{Internal nonlinear interactions} which couple three waves to each other.  These interactions enable waves that are linearly excited by the tide to transfer energy to waves that are not linearly excited.  This nonlinear coupling includes both resonant and non-resonant nonlinear interactions (\\S\\S~\\ref{sec:nonlinear_tide}, \\ref{sec:stability_analysis}, \\& \\ref{sec:inhomog_driving}).  If the linearly excited waves systematically lose energy to other waves via nonlinear coupling, then the nonlinear interactions can substantially modify the orbital/rotational evolution relative to that predicted by linear theory. (2) \\emph{External nonlinear interactions} which resonantly couple two internal waves directly to the tide through their gravitational multipole moments (\\S~\\ref{sec:simple_nltide}). These interactions allow waves that are not linearly resonant with the tide to be directly driven by the tidal potential. The damping of these nonlinearly driven waves can also act as a source of enhanced dissipation relative to linear theory.  We have presented two related approaches to solving for the nonlinear tidal response of a star.  In the first, we take the background state of the star to be spherical and unperturbed by the companion and we expand the Lagrangian displacement and velocity associated with the tide (both linear and nonlinear) as a sum over eigenmodes, with each mode weighed by its amplitudes $q_a$ (\\S~\\ref{sec:method1}).  This extends one of the standard methods for studying linear tides to include nonlinear interactions.  In the second method, we instead take the background state of the star to include the linear tidal solution and we expand the nonlinear correction to the tidal solution as a sum over eigenmodes, with each mode weighed by a different mode amplitude $r_a$ (\\S \\ref{sec:method2}).  These two methods are formally equivalent but we find that each is useful for understanding different aspects of tides in stars.  The second method of solving explicitly for the nonlinear correction to the tidal response of a star is particularly useful for understanding the nonlinear stability of standard linear tidal theory (\\S~\\ref{sec:stability_analysis}).  Since we focus on slowly rotating stars, the stellar modes included in our treatment are p-modes and g-modes; the latter are particularly important because they can have periods comparable to that of the binary system.  The orthogonality and completeness of the stellar eigenmodes enables us to convert the nonlinear partial differential fluid equations into a coupled network of ordinary differential equations describing the evolution of each mode amplitude.  We now summarize the physics contained in these equations, focusing for concreteness on the method of expanding the full tidal solution (linear and nonlinear) as a sum over eigenmodes.  In this case, the resulting amplitude equations have linear terms and nonlinear terms. The linear terms are standard (e.g., \\citealt{Press:77}) and include driving by the tidal force and damping by radiative diffusion.\\footnote{ While dissipation due to eddy viscosity in convection zones is important for the long wavelength equilibrium tide, we ignore it in this paper, focusing instead on thermal diffusion damping in the radiative zone, which is more important for short wavelength, low frequency g-modes in the radiative core.} This is the basic physics included in most previous studies of tides, with one addition: by solving the mode amplitude equations, our formalism allows the modes to be dynamic in their interaction with the orbit, rather than assuming harmonic response at the forcing frequency.  The nonlinear terms in the equations of motion are parameterized by the coefficients $\\kappa_{abc}$ and $U_{ab}$ (see eqs. [\\ref{eq:kappa}] and [\\ref{eq:Uab}] in \\S~\\ref{sec:eom}); these describe the internal (three wave) and external (nonlinear tidal driving) nonlinear interactions, respectively.  Accurately calculating these coupling coefficients is nontrivial -- these technical details are described in Appendix \\ref{sec:app:coef_in_amp_eqn}, where we pay considerable effort to ensuring that the coupling coefficients can be accurately calculated numerically.  Three wave coupling has been considered extensively in the literature, both in stellar seismology (e.g., \\citealt{Dziembowski:82,Wu:01}) and more specifically in the context of tidally excited oscillations in binary systems (e.g., \\citealt{Kumar:96}).  Nonlinear tidal driving has not, to our knowledge, been studied in any detail before (\\citealt{Papaloizou:81} considered nonlinear tidal driving but neglected three-wave coupling which we found cancels strongly with nonlinear tidal driving).  Both types of interactions can lead to a variety of physical effects including the non-resonant excitation of higher frequency modes and the resonant excitation of lower frequency modes (parametric instability).We have focused our analysis on determining (1) the conditions under which the linear tidal flow is unstable to the parametric instability (\\S\\S~\\ref{sec:dyntide} \\& \\ref{sec:eqtide}) and (2) the efficiency of nonresonant excitation of g-modes and p-modes by nonlinear coupling to the linear tidal flow (\\S~\\ref{sec:inhomog_driving}).   In the course of computing nonlinear coupling coefficients and growth rates for solar-type stars, we found the {\\em a priori} surprising result that the external nonlinear driving of g-modes ($U_{ab}$) is almost completely canceled by a portion of the three wave coupling to the linear tide ($\\kappa_{abc}$).  The implication is that these two effects are intimately related. The nature of this relationship can be most easily appreciated if we explicitly solve for the nonlinear correction to the linear tidal solution (method 2 in \\S~ \\ref{sec:second_order_eom}) and consider the linear flow as a superposition of an ``equilibrium tide\" and a ``dynamical tide\"; the former is the nearly hydrostatic part of the linear tidal response and the latter is the wave-like, resonant part of the response. The three wave coupling coefficient that describes the nonlinear driving of a pair of `daughter' waves then has a contribution from the equilibrium tide $\\kappa_{ab}^{(\\rm eq)}$ (eq. [\\ref{eq:kappa_eq}]) and the dynamical tide $\\kappa_{ab}^{(\\rm dyn)}$ (eq. [\\ref{eq:kappa_dyn}]). Since the equilibrium tide contains the vast majority of the tidal energy, one might expect it to be more prone to nonlinear instability than the dynamical tide (e.g., \\citealt{Press:75,Kumar:98a}).  In fact, we find the opposite: the daughters internal driving via three wave coupling to the equilibrium tide nearly cancels with their external nonlinear driving by the tidal potential ($\\propto U_{ab}$). As a result, the effective coupling of daughters to the equilibrium tide is much weaker than their coupling to the dynamical tide, i.e., $|\\kappa_{ab}^{(\\rm eq)}| \\ll |\\kappa_{ab}^{(\\rm dyn)}|$ (see \\S~\\ref{sec:driving_rate} and Appendices \\ref{sec:app:alternative_eom} and \\ref{sec:properties_of_coefficient}).  Physically, this is because internal nonlinear driving by the equilibrium tide and nonlinear tidal driving are fundamentally part of the same process; together they describe the nonlinear excitation of daughter modes by the nearly hydrostatic response of the star to its companion.  Another reason for the weak nonlinear driving by the equilibrium tide in solar-type stars is that very little of the energy of the equilibrium tide resides in the core of the star where the low frequency g-modes that the equilibrium tide attempts to drive have most of their energy. At a technical level, the subtlety of calculating the equilibrium tide driving correctly highlights the importance of the detailed calculation presented in this paper.  For example, an order of magnitude estimate of the nonlinear driving by the shear of the equilibrium tide ($\\kappa_{ab}^{(\\rm eq)} \\sim 1$) overestimates the true driving by almost two orders of magnitude (\\S~\\ref{sec:eq_tide_driving_rate} and Appendix \\ref{sec:Jeq}; Figs. \\ref{fig:Jplus2kap} and \\ref{fig:Jplus2kap_dn}).  The equilibrium tide driving of daughter modes, when present, is relatively evenly spread throughout the radiative zone (left panel of Fig. \\ref{fig:Jplus2kap_dn}); by contrast, the driving by the dynamical tide is highly concentrated in the core of the star (Fig. \\ref{fig:kappa}).  Hence the two mechanisms for seeding fluid instability are physically quite different.  Quantitatively, we find that the equilibrium tide is only unstable to nonlinear driving of resonant g-modes in solar binaries if the orbital period is $\\lesssim 2-5$ days (\\S~\\ref{sec:eqtide}).  This implies that instability of the equilibrium tide in slowly rotating solar-type stars cannot help explain the observed circularization of solar binaries out to orbital periods of $\\simeq 10-15$ days. An interesting remaining possibility is that the equilibrium tide in {\\em rotating} solar-type stars may efficiently couple nonlinearly to inertial waves.  In particular, inertial waves in solar type stars have most of their energy in the outer convection zone where much of the energy in the equilibrium tide also resides.  Recently, the elliptical instability has been invoked as a source of nonlinear driving in close binary systems (\\citealt{Lebars:10, Cebron:11}; see also \\citealt{Press:75, Seguin:76}). The elliptical instability is clearly related to the equilibrium tide instability considered here; both instabilities involve the destabilization of internal modes of oscillation by the long wavelength, hydrostatic, tidal perturbation. They are not equivalent, however, as the elliptical instability drives inertial modes whereas we focus on g-modes.  For solar type stars, we find that the dynamical tide driving rates of short wavelength g-modes ($\\propto \\kappa_{ab}^{(\\rm dyn)}$) are orders of magnitude larger than the equilibrium tide driving rates. In order to determine when linear theory is valid, we computed the parametric instability growth rates over a range of orbital periods and companion masses (\\S~\\ref{sec:dyntide}).  We find that for orbital periods from days to weeks, the dynamical tide part of the linear tidal flow is unstable, even for companion masses much smaller than a Jupiter mass (see Fig. \\ref{fig:massperiod}).  The linear tidal solution in the star is invalid even for a 10 Earth mass planet if the orbital period is $\\lesssim 5$ days! The degree to which the true tidal flow differs from the linear one depends on the saturation of these nonlinear instabilities, which will be investigated in a future paper.  This question is of particular current interest given the large number of low mass planets orbiting solar-type stars being discovered by transit surveys such as the \\emph{Kepler} telescope \\citep{Borucki:11}.  The orbital properties of these systems are likely shaped in part by the nonlinear instabilities discussed in this paper.  In our study of the parametric instability of the dynamical tide we found that it is subject to a ``collective'' version of the parametric instability. In the literature, parametric growth rates have traditionally been derived by considering how a single pair ($N\\le 2$) of coupled short wavelength daughters interact with the time-dependent background of a parent mode (e.g., \\citealt{Dziembowski:82}). We have extended this analysis to allow for $N\\gg2$ daughters, each of which is coupled to the parent and the other daughters.  We find that for solar type stars, large sets of daughters are coherently driven by the dynamical tide ($N\\approx 10^2-10^3$ for periods $P\\la10\\trm{ days}$ and companions $M'/M\\ga few\\times10^{-4}$). When collectively unstable, the daughters all grow as a single coherent unit, i.e., maintaining phase coherence. This coherence is self-consistently generated by the nonlinear interactions even if it is not present in the initial amplitudes/phases of the modes (Fig. \\ref{fig:simple_3wave_collective}).  In the collective parametric instability, the growth rates of the daughters are significantly larger than in standard three wave coupling, by a factor of $\\approx N/2$ (Fig. \\ref{fig:growth_collective}).  As a result, the collective instability may have important implications for the nonlinear damping of the dynamical tide. More generally, it will also be important to revisit previous astrophysical applications of parametric instability to understand the consequences of the restrictive focus on three-mode coupling, which does not necessarily capture the parametric instability threshold or fastest growth rate.   In their study of the dynamical tide in solar-type binaries, \\citet{Goodman:98} noted that short-wavelength gravity waves traveling inwards from the radiative-convective boundary (where they are excited) reach sufficiently large amplitudes that the waves ``break'' at the center of the star.  \\citet{Barker:10} and \\citet{Barker:11} confirmed this result numerically.  For hot Jupiter companions with masses $\\lesssim 2 M_{\\rm Jupiter}$, however, they found that the incoming gravity waves do not break in the core of the star (see also \\citealt{Ogilvie:07}).  In an effort to understand what happens in the regime where the incoming wave does not break in the core, \\citet{Barker:11b} performed a stability analysis of a non-linear standing internal gravity wave in the central regions of a cylindrical ``star\". Although their methods differ from ours, they obtain (non-collective) daughter driving rates similar to ours (see \\S~\\ref{sec:dyn_tide_driving_rate}). Like us, \\citet{Barker:11b} do not attempt to solve for the saturation of the unstable daughters and thus do not explicitly solve for the tidal $Q$ of the star. However, they argue that the driving rate of the fastest growing daughter mode places a lower bound of $Q\\ga 10^7$. We are not convinced that this result is correct for two reasons. First, both our study and theirs find that the parent excites many daughter pairs, many of which have (non-collective) growth rates just slightly below that of the fastest growing mode. These excited daughters will each  drain energy from the parent. Since there are so many such rapidly growing daughters, it seems possible that the parent will lose more energy to all of these modes than it does to the single fastest growing mode. Secondly, \\citet{Barker:11b} do not seem to observe the collective instability, and thus might be underestimating the rate at which daughters drain energy from the parent.\\footnote{In the absence of collective driving, the driving rate $\\Gamma_{bb}\\propto \\omega/n$ (see \\S~\\ref{sec:dyn_tide_driving_rate}). Collective driving increases the driving rate of each daughter by a factor of approximately $N\\approx n$ and therefore the collective driving rate $\\Gamma_{bb}\\propto \\omega$. \\citet{Barker:11b} find $\\Gamma_{bb}\\propto \\omega/n$ even above the collectively instability threshold.}  One possible explanation for why \\citet{Barker:11b} do not see the collective instability is that their diffusivity, which they chose in order to obtain numerically converged solutions, is set so high that it artificially suppresses collective driving.  In \\S~\\ref{sec:dyntide} we show that the companion (planet) mass above which the collective instability sets in lies well above the threshold for the standard three-mode instability (but still well below the wave breaking threshold; see Fig. \\ref{fig:massperiod}). Thus, if their diffusivity is set too high compared to the true value in a solar-type star, they might be in the regime where standard driving is unstable but collective driving is artificially stabilized.   In addition to the driving of resonant high $\\ell$ daughter waves, there is efficient non-resonant three mode coupling (nonlinear inhomogeneous driving) between the dynamical tide, the equilibrium tide,  and a wide range of stellar p-modes and g-modes (\\S~\\ref{sec:inhomog_driving} and Figs.~\\ref{fig:EinhomogP} \\& \\ref{fig:linlin_coef}).  How this coupling saturates is, however, unclear.  Even if these alternative nonlinear couplings turn out to be more important than wave breaking in damping the dynamical tide, they are unlikely to modify the orbital evolution due to the dynamical tide calculated by \\citet{Goodman:98}; this is because the linear energy input rate into the dynamical tide is independent of its damping rate in the limit of efficient damping.  Nonetheless, identifying the correct damping mechanism for the dynamical tide in solar binaries is important because it determines where in the star the tidal energy is dissipated.  In some cases, we find that the parametric growth times due to driving by the dynamical tide can be so short that the excited daughter waves do not have time to travel through the core of the star (at the group velocity) before their amplitudes are expected to become highly nonlinear (\\S~\\ref{sec:global_vs_local}).  This is particularly true for the ``collective'' parametric instability (Fig. 10).  In this case, the standing wave approximation employed here may become inaccurate and a traveling wave point of view may be more appropriate. While this does not alter our conclusions about the stability of the linear solution, it will likely affect how the nonlinearities saturate and thus the rate of tidal dissipation.  To conclude, we note that the general formalism developed in this paper can be extended to the tidal forcing of other types of stars. For solar-type stars, the resonant g-mode amplitudes are large at the center of the star, where the equilibrium tide amplitude is small. In other stars, resonant wave amplitudes may be large at the stellar surface, where the equilibrium tide is also large. Equilibrium tide driving may thus play a more prominent role in those stars."
    },
    "1107/1107.0494_arXiv.txt": {
        "abstract": "{ Enhancement of nuclear pasta formation due to multi-nucleus simultaneous collision is presented based on time-dependent density functional calculations with periodic boundary condition. This calculation corresponds to the situation with density lower than the known low-density existence limit of the nuclear pasta phase. In order to evaluate the contribution from three-nucleus simultaneous collisions inside the cosmic matter, the possibility of multi-nucleus simultaneous collisions is examined by a systematic Monte-Carlo calculation, and the mean free path of a nucleus is obtained. Consequently the low-density existence limit of the nuclear pasta phase is formed to be lower than believed up to now. } % ",
        "introduction": "Nuclear pasta, whose existence and role were suggested in association with the collapse driven supernova explosions \\cite{watanabe,sowa}, is of great importance with respect to the astrophysical synthesis of chemical elements. Indeed nuclear pasta is known to change the transparency rate of neutrinos \\cite{ravenhall,sato}, which play a role of trigger in supernova explosions giving rise to a wide variety of chemical elements. Therefore the existence of nuclear pasta at low densities is expected to have a great impact on clarifying the historical evolution of chemical elements and also on revising the low-density and low-temperature part of the nuclear phase diagram \\cite{sonoda,lamb}. In particular, here we are concerned with its existence in the outermost part of protoneutron stars (for the review of supernova mechanisms, see \\cite{bethe}).  In this paper, based on three-dimensional time-dependent density functional calculations, a dynamical formation of nuclear pasta in the low-density and low-temperature situations is demonstrated by 3-nucleus simultaneous heavy-ion collision. This calculation simulates the situation outside of the known low-density existence limit of nuclear pasta phase. In order to examine the validity of this calculation, the possibility of multi-nucleus simultaneous collisions is studied by a systematic three-dimensional Monte-Carlo calculation, and the mean free path of a nucleus at a given density and temperature is obtained. Note that for this kind of nuclear pasta formation, the presence or absence of the fast charge equilibration process \\cite{iwata} determines the upper boundary of nuclear pasta phase with respect to the temperature \\cite{iwata2}. ",
        "conclusions": "Enhancement of nuclear pasta formation due to 3-nucleus simultaneous collisions has been presented based on time-dependent density functional calculations (using SLy6 and SKI3 force parameter sets). Multi-nucleus simultaneous collisions between more than two nuclei have been confirmed to occur in cosmic matter with density higher than $0.01 \\rho_0$. Therefore, nuclear pasta formation is possible in the outermost part of protoneutron stars, and neutrino radiation is not as easy as expected. Eventually, the low-density existence limit of nuclear pasta phase is suggested to be lower than believed up to now. \\\\   The author thanks to Profs. K. Iida, N. Itagaki, J. A. Maruhn and T. Otsuka. This work was supported by the Helmholtz Alliance HA216/EMMI."
    },
    "1107/1107.2943_arXiv.txt": {
        "abstract": "We observed the Sh~2-233IR (S233IR) region with better sensitivity in near-infrared than previous studies for this region.  By applying statistical subtraction of the background stars, we identified member sources and derived the age and mass of three distinguishable sub-groups in this region: Sh~2-233IR NE,  Sh~2-233IR SW, and the ``distributed stars'' over the whole cloud.  Star formation may be occurring sequentially with a relatively small age difference ($\\sim0.2-0.3$ Myrs) between subclusters.   We found that the slopes for initial mass function ($\\Gamma\\sim -0.5$) of two subclusters are flatter than that of Salpeter, which suggests that more massive stars were preferentially formed in those clusters compared to other Galactic star-forming regions.  These subclusters may not result from the overall collapse of the whole cloud, but have formed by triggering before the previous star formation activities disturbed the natal molecular cloud.   Additionally, high star formation efficiency ($\\gtrsim$40\\%) of the subclusters may also suggest that stars form very efficiently in the center of NE. ",
        "introduction": "Within 2 kpc, $70-90$\\% of stars have been found to form in the clusters \\citep{lada03}.   These newly formed stars, however, are embedded deeply in dense molecular cloud and there exists various difficulties in identifying sources and deriving physical parameters.  Several statistical tools, such as the K-band luminosity function (KLF), the color-color or color-magnitude diagrams, etc., have been successfully used to constrain the characteristics of deeply embedded stellar clusters \\citep[e.g.,][]{lada03}.   Development of new sensitive equipments, such as, large mosaic IR array, enables us to observe embedded stellar cluster in greater detail than previous effort.  The Sh~2-233IR region is a well studied star-forming region in our Galaxy \\citep[e.g.,][]{porras00} and it provides an ideal laboratory for understanding properties of embedded clusters.  Therefore, we revisit this region with higher sensitivity near-infrared data taken toward a wider region to better constrain the properties of the embedded stellar cluster than previous studies.  Toward the direction of the Galactic Anticenter, Sh~2-233IR (hereafter S233IR), as a part of the Sh~2-235 GMC complex \\citep{ry08}, is located at a distance of about 1.8 kpc \\citep{porras00}, in association with four extended HII regions, Sh~2-231, 232, 233, and 235 \\citep{heyer96}.  Its position coincides with an IRAS source, IRAS 05358+3543 ($\\alpha_{2000}=05^h39^m10^s$ $\\delta_{2000}=+35^{\\circ}45^\\prime19^{\\prime\\prime}$).  S233IR is classified as a massive star formation region \\citep{sri02}, showing CO outflows \\citep{snell90} and various maser emissions associated with this region  \\citep{henning92,tofani95,beuther02b,menten91,minier05}.    The $K^\\prime$ band image of this region shows many bright stellar sources and also extended nebulous features associated with dust emission \\citep{hodapp94}.  The two embedded young clusters, Sh~2-233IR SW (hereafter SW, located in the south-west direction from the center) and Sh~2-233IR NE (hereafter NE, in the north-east direction), are notable in this region, with remarkable H$_2$ bow shocks associated with the NE cluster \\citep{porras00}. In addition, numerous studies have been reported, especially for the NE cluster, including polarimetric observations \\citep{jiang01,yao00},  molecular outflows \\citep{beuther07,mz04,beuther02a,cesaroni99,lari99}, and mid-infrared sources \\citep{longmore06}.  In this {\\it paper}, we revisit this S233IR region with wider field of view, higher resolution, and better sensitivity data in near-infrared and radio wavelengths.   Our goal is to understand star formation history in terms of age, star formation efficiency and initial mass function.  We summarize our observations at radio and near-infrared in \\S~\\ref{obs}, and the observed results in \\S~\\ref{results}. Discussion is given in \\S~\\ref{discuss}, and the summary in \\S~\\ref{summary}.  This is the first paper in the series of our work for the S233IR and associated region.  Here we focus on properties of the embedded stellar clusters in the S233IR region.  The H$_2$ shock features and cloud kinematics have been studied by \\cite{adam09}, and the extended nebulous $K$-band features, young stellar object and star formation in a larger area will be discussed in the second paper \\citep{chy09b}. ",
        "conclusions": "\\label{discuss}  Based on the radio and extinction map \\citep{heyer96,ry08}, the S233IR molecular cloud is a part of the much larger Auriga GMC complex, which is  associated with several large HII regions and active star-forming regions.  The large scale ($\\sim$ a few degrees) morphology of this GMC complex looks very clumpy and is composed of several HII regions, such as Sh~2-231, 232, 233 and 235 \\citep{ry08}.   The CO emission shows a filamentary and arc-like structure associated with three HII regions.   At the northeast direction, an HII region,  Sh~2-231, is $\\sim10\\arcmin$ away from S233IR  \\citep{mz04}.  In addition, there are many newly formed OB stars, such as those in Aur OB1, associated with this GMC complex \\citep[][Figure 1]{ry08}.  The S233IR clump appears to be the latest of star formation events.  Three groups of stars were identified in S233IR region, SW, NE and the distributed stars.  The background contamination in distributed stars were eliminated by a statistical method.  The age differences among these subgroups were first noticed by \\citet{porras00} using the $J$-band analysis.  Based on better data quality of our work, we derived ages $\\sim$0.5 Myrs for SW, $\\sim$0.3 Myrs for NE, and $\\sim$1.5 Myrs for the distributed stars.   The distributed stars could be the first generation of stars formed in this cloud.  The SW and NE clusters, with the relative distance of about 0.5 pc at a 1.8 kpc distance, are later generations.   The age differences between stellar groups suggests sequential star formation, which has also been found in other star-forming regions, such as NGC~7538 \\citep{ojha04}, Sh~2-247 \\citep{puga08}, Sh~2-157 \\citep{chen09} and NGC~3576\\citep{purcell09}.  It is interesting that the S233IR cloud, a typical star-forming cloud in size ($r\\sim$1.5 pc) and mass ($\\sim1100$ M$_\\odot$), contains sub-clusters formed in different epochs and star formation still continues without disrupting its natal molecular cloud.  Stars appear to be forming continuously in the youngest NE cluster.   For the distributed stars, the derived IMF slope,  $\\Gamma\\sim -0.5$, is flatter than in other  star-forming regions in our Galaxy,   which indicates that more massive stars have been formed preferentially in this cloud than in other typical star-forming regions.   The $\\Gamma$ value of the NE cluster, the latest star-forming sub-core in this cloud, is even flatter than the distributed stars,  suggesting that the probability of massive star formation is higher at this region.   This result is consistent with the radio observation toward the center of NE cluster \\citep{beuther02a}.   Eventually the cluster will disperse the natal molecular cloud and appear as visible stars.  Once the massive stars are formed, the radiation from massive star will clear the molecular could and stop the star formation process in the core region.    The enhancement of  the local SFE is also found in the central dense core as we derived in \\S\\ref{sfe}, although  we have to take various uncertainties into consideration, for example, the uncertain total H$_2$ density, the cloud boundary, member star determination, undetected sources especially near the central part of the sub-clusters showing saturated infrared emission, etc.  The local high SFE ( $\\gtrsim$40\\%) suggests that the later generation of stars can be formed very efficiently in dense cores in the same molecular cloud.    In the outskirt of the cloud, where the first stellar generation is located, the SFE is consistent with a typical molecular cloud.    The trigged star formation may have enhanced the formation of massive stars, which  results in a high SFE and flatter IMF slope in the central region.  It is known that the NE cluster is the latest star forming core in S233IR, which is associated with the powerful energetic features, such as shocked H$_2$ gas flows \\citep{adam09}.  These energetic activities are thought to be a part of massive star formation process.   The S233IR region provides an interesting example of active and energetic massive star formation, which could be further studied with next generation instruments."
    },
    "1107/1107.2202_arXiv.txt": {
        "abstract": "{In addition to coherent pulsation, many accreting neutron stars exhibit flaring activity and strong aperiodic variability on time scales comparable to or shorter than their pulsation period. Such a behavior shows that the accretion flow in the vicinity of the accretor must be highly non-stationary. Observational study of this phenomenon is often problematic as it requires very high statistics of X-ray data and a specific analysis technique.} {In our research we used high-resolution data taken with \\textsl{RXTE} and \\textsl{INTEGRAL} on a sample of bright transient and persistent pulsars, to perform an in-depth study of their variability on time scales comparable to the pulsation period -- ''pulse-to-pulse variability''.} {The high-quality data allowed us to collect individual pulses of different amplitude and explore their X-ray spectrum as a function of pulse amplitude. The described approach allowed us for the first time to study the luminosity-dependence of pulsars' X-ray spectra in observations where the averaged (over many pulse cycles) luminosity of the source remains constant.} {In all studied pulsars we revealed significant spectral changes as a function of the pulse height both in the continuum and in the cyclotron absorption features. The sources appear to form two groups showing different dependencies of the spectrum on pulse height. We interpret such a division as a manifestation of two distinct accretion regimes that are at work in different pulsars.} {} ",
        "introduction": "Strong aperiodic variability of the X-ray flux is a well known characteristic of many X-ray binary systems. First discovered in black hole candidates (with Cyg~X-1 being the most remarkable case), it was later shown to be a common feature also among accreting neutron stars \\citep[see e.g.][for a review]{Belloni:Hasinger:90}. Although \\emph{homogeneity} and \\emph{stationarity} of the accretion flow is often assumed in calculations dealing with physics inside the X-ray emitting structure on a neutron star (as it greatly simplifies the mathematical treatment of the problem), the observed variability clearly indicates that the accretion flow in the close vicinity of accretor must be highly non-stationary. On the other hand, the penetration of matter into the magnetosphere of the neutron star around the stopping radius is believed to be governed by various kinds of instabilities which plasma supported by magnetic field is subject to \\citep[see e.g.][]{Ghosh:Lamb:79,Kulkarni:Romanova:08,Bozzo:etal:08}. The instabilities will naturally lead to fragmentation of a continuous flow into more or less isolated blobs or filaments. More focused theoretical studies of the non-stationary accretion problem have been performed e.g. by \\citet{Morfill:etal:84,Demmel:etal:90, Orlandini:Boldt:93} who have explicitly shown that the inhomogeneity of the flow is generally expected in accreting pulsars and that it does not only arise from the original inhomogeneity of the accreted matter (e.g. ''clumps'' in the stellar wind or variations of matter supply from the donor star) but can naturally be produced by instabilities close to the magnetospheric boundary of the neutron star.  So far, the aperiodic variability of accreting pulsars has mostly been studied by means of power spectra of their high time-resolution light curves \\citep[e.g.][]{Belloni:Hasinger:90,Pottschmidt:etal:98,Revnivtsev:etal:09}. In our work, however, we concentrate on the variability of individual X-ray pulses often referred to as \\emph{pulse-to-pulse variability}. Such kind of variability seems to be a common phenomenon among accreting pulsars and has been reported by several authors for bright outbursts of some transient or strongly variable sources: \\citet{Frontera:etal:85} for A\\,0535+26, \\citet{Staubert:etal:80,Kretschmar:etal:00} for Vela\\,X-1, \\citet{Tsygankov:etal:07} for 4U\\,0115+63. However, a detailed study of the pulse-to-pulse variability is usually limited by the photon statistics, especially for relatively fast pulsars (with a spin period of a few seconds or less). In the present work, we used the publicly available archival data from high-sensitivity X-ray detectors onboard \\textsl{RXTE} and \\textsl{INTEGRAL} taken on a set of four bright accreting pulsars -- V\\,0332+53, 4U\\,0115+63, A\\,0535+26, and Her\\,X-1.  The four accreting pulsars are well established cyclotron line sources, i.e. their spectra contain \\emph{Cyclotron Resonant Scattering Features (CRSFs)}. Such features appear as absorption lines due to resonant scattering of photons off the relativistic plasma electrons at Landau levels \\citep[see e.g.][]{Truemper:etal:78,Isenberg:etal:98,Araya:Harding:00}. The CRSFs, if detected, provide a direct way to measure the magnetic field strength at the site of X-ray emission as the energy of the fundamental line and the spacing between the harmonics are proportional to the $B$-field strength. In some sources the cyclotron line energy was found to change with luminosity. In V\\,0332+53 and 4U\\,0115+63 a negative correlation of the line energy with luminosity was observed \\citep{Tsygankov:etal:07,Tsygankov:etal:10,Mowlavi:etal:06,Mihara:etal:04} while for Her~X-1, \\citet{Staubert:etal:07} reported a positive correlation of the line energy with luminosity.  A description of the observational data is provided in Sect.\\,\\ref{sec:obs}. Using a special analysis technique (described in Sect.\\,\\ref{sec:techn}), we collected individual pulses of different amplitudes and studied the differences in their X-ray spectra. As a result, we revealed significant correlations of different spectral parameters with pulse amplitude (Sect.\\,\\ref{sec:results}). Based on the sign of the correlations the explored pulsars can be divided into two groups. We argue that the two groups correspond to two distinct regimes of accretion that are at work in different sources: local sub- and super-Eddington regime (see discussion in Sect.\\,\\ref{sec:discussion}). The summary and conclusions are provided in Sect.\\,\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In the presented work we studied the pulse-amplitude-resolved spectral variability in a sample of four bright accreting pulsars using the high-resolution X-ray data taken with \\textit{RXTE} and \\textit{INTEGRAL}. Our analysis revealed for the first time the spectral differences between X-ray pulses of different amplitudes, both in the spectral continuum and in the resonant cyclotron feature.  For the pulsars where a negative correlation of the cyclotron line energy with flux on longer time scales has been reported previously we found a similar (negative) correlation on the time scale of singe pulse cycles. In Her~X-1 we found a positive correlation of the line energy with pulse height which is consistent with the long-term correlation of the CRSF energy and the maximum main-on flux reported by \\citet{Staubert:etal:07}. Our analysis revealed a positive correlation of the cyclotron line energy with pulse amplitude in A\\,0535+26 where no correlation of the CRSF energy with flux has so far been reported.  The pulsars in our sample show two different types of spectral dependencies on pulse amplitude. We argue that the different behaviors reflect two distinct accretion regimes (local sub- and super-Eddington) realized in different pulsars, as it was previously proposed to explain the peculiarity of Her~X-1 compared to higher-luminosity pulsars."
    },
    "1107/1107.2399_arXiv.txt": {
        "abstract": "{} {In the last four decades it has been observed that solar flares show quasi-periodic pulsations (QPPs) from the lowest, i.e. radio, to the highest, i.e. gamma-ray, part of the electromagnetic spectrum. To this day, it is still unclear which mechanism creates such QPPs. In this paper, we analyze four bright solar flares which show compelling signatures of quasi-periodic behavior and were observed with the Gamma-Ray Burst Monitor (\\gbm) onboard the \\fermi satellite. Because \\gbm covers over 3 decades in energy (8~keV to 40~MeV) it can be a key instrument to understand the physical processes which drive solar flares.} {We tested for periodicity in the time series of the solar flares observed by \\gbm by applying a classical periodogram analysis. However, contrary to previous authors, we did not detrend the raw light curve before creating the power spectral density spectrum (PSD). To assess the significance of the frequencies we made use of a method which is commonly applied for X-ray binaries and Seyfert galaxies. This technique takes into account the underlying continuum of the PSD which for all of these sources has a $P(f)\\sim f^{-\\alpha}$ dependence and is typically labeled red-noise.} {We checked the reliability of this technique by applying it to a solar flare which was observed by the Reuven Ramaty High-Energy Solar Spectroscopic Imager (\\rhessi) which contains, besides any potential periodicity from the Sun, a 4~s rotational period due to the rotation of the spacecraft around its axis. While we do not find an intrinsic solar quasi-periodic pulsation we do reproduce the instrumental periodicity. Moreover, with the method adopted here, we do not detect significant QPPs in the four bright solar flares observed by \\gbm. We stress that for the purpose of such kind of analyses it is of uttermost importance to appropriately account for the red-noise component in the PSD of these astrophysical sources.} {} ",
        "introduction": "Over the last 40 years quasi-periodic pulsations (QPP) in solar flares have been reported from observations across the electromagnetic spectrum, i.e. from radio waves up to the high energetic gamma-rays ranging from sub-second timescales up to several minutes  \\citep[e.g.][]{parks69, ofsu06, ligan08, nakmel09, nak10}. While there seems to be an overwhelming amount of observational data, the underlying physical mechanism which could generate such QPPs remains still a mystery. Possible processes that are considered are modulation of electron dynamics by magnetohydrodynamic (MHD) oscillations \\citep{zait82}, periodic triggering of energy releases by MHD waves \\citep{foull05, nakfoull06}, MHD flow overstabilities \\citep{ofsu06} and oscillatory regimes of magnetic reconnection \\citep{kliem00}.  In this paper we will present time series and periodogram analyses of four solar flares which show an intriguing quasi-periodic behavior in their light curves. All of these solar flares were observed by the \\fermi Gamma-Ray Burst Monitor \\citep[GBM,][]{meegan09}. \\gbm is one of the instruments onboard the \\fermi Gamma-Ray Space Telescope \\citep{atwood09} launched on June 11, 2008. Specifically designed for gamma-ray burst (GRB) studies, \\gbm observes the whole unocculted sky with a total of 12 thallium-activated sodium iodide (NaI(Tl)) scintillation detectors covering the energy range from 8~keV to 1~MeV and two bismuth germanate scintillation detectors (BGO) sensitive to energies between 150~keV and 40~MeV \\citep{meegan09}. Therefore, \\gbm offers superb capabilities for the analyses of not only GRBs but solar flares as well.  The analysis and interpretation of power spectral density (PSD) of solar flares is, in general, difficult. A variety of astrophysical sources (such as X-ray binaries, Seyfert galaxies \\citep[e.g.][]{lawrence87, mark03} and GRBs \\citep{ukwatta09, cenko10}), show erratic, aperiodic brightness changes. Solar flares exhibit similar aperiodic variations with the general time profile being a sharp impulsive phase followed by a slower decay phase. Solar flares, together with many other astrophysical sources, thus have very steep power spectra in the low-frequency region. This type of variability is known as red-noise \\citep[e.g.][]{groth75, deeter82,vaughan05}. When determining the significance of possible periodicities in the PSD the red-noise has to be accounted for in order not to severely overestimate the significance of identified frequencies \\citep{lachgup09}. In this paper we will account for the red-noise properties when performing a periodogram analysis.  This paper is organized as follows. In Sect.\\ref{sec:rednoise} we briefly present the methodology of the time-series analysis. We provide an overview of the red-noise properties in astrophysical sources and demonstrate the importance of the red-noise when estimating significances. In Sect.\\ref{sec:gbm} we will present light curves and periodograms of solar flares which were observed by \\gbm. Finally, in Sect.\\ref{sec:conc}, we will summarize and conclude. ",
        "conclusions": "\\label{sec:conc} We have analyzed light curves of five solar flares observed by both \\gbm and \\rhessi. We tested the data for the presence and significance of QPPs accounting for the overall shape of the PSD by applying the method introduced by \\citet{vaughan05}. First of all, this technique was validated and tested by applying it on raw data of the \\rhessi satellite. Any \\rhessi light curve has an inherent period caused by the rotation of the space craft around its axis. With the method adopted here we successfully retrieve this well known 4~s period. However, we were not able to confirm the previously reported QPP of 40~s in the very same solar flare \\citep{nak10}. An additional check was performed by applying the method to SPI-ACS data of the giant flare of the well known SGR~1806-20. This magnetar has a rotation period of $7.56$~s which, together with several harmonics, could be recovered unambiguously with the procedure presented here. These two tests gave us confidence that the method is appropriate to test the significance of QPPs in red-noise dominated solar flare time series.  The routine was then applied to four solar flares observed by \\gbm. Although all of these solar flares showed very promising quasi-periodic features in their detrended light curves at first none of these were significant ($>$ 3~$\\sigma$). Previous authors \\citep[e.g.][]{inglis08,inglis09,zim10,nak10} who claim a significant QPP detection, applied a standard Lomb-Scargle analysis to detrended solar flare light curves. Our investigation of this method suggests the power in the low frequency range is being artificially suppressed, which can lead to misleading values for the significances of features in the PSD. A periodogram analysis should always be performed on the raw and undetrended light curve as was done here.  We stress once more that not only solar flares but many astrophysical sources (X-ray binaries, Seyfert galaxies, GRBs) suffer from steep power spectra in the low-frequency range. Such power spectra make a periodogram analysis not trivial, the interpretation of a peak in a PSD more complex and an estimation of its significance an important issue. Again we emphasize that red-noise is an intrinsic source property. In other words, having shown that the variations in the corresponding solar flares are not quasi-periodic at the 3$\\sigma$ level does not mean that these variations are not real. The corresponding flux changes are sometimes dramatic, reaching a factor of a few within a few tens of seconds. Also, these variations occur in phase at different X-ray to gamma-ray energies, and other flares have been observed to also occur in phase with microwave radio emission \\citep[e.g.][]{nak10,fou10}.  While we think that it is not necessary to invoke oscillatory regimes of plasma instabilities \\citep{nakingl10, rez11}, we note that the physics of these variations is certainly interesting and worth further studies."
    },
    "1107/1107.5570_arXiv.txt": {
        "abstract": "We combine data from the Australia Telescope National Facility and the Swedish ESO Submillimeter Telescope to investigate the neutral interstellar medium (ISM) in AM0644-741, a large and robustly star-forming ring galaxy. The galaxy's ISM is concentrated in the $42$ kpc diameter starburst ring, but appears dominated by atomic gas, with a global molecular fraction ($f_{\\rm mol}$) of only $0.079 \\pm 0.005$. Apart from the starburst peak, the gas ring appears stable against the growth of gravitational instabilities ($Q_{\\rm gas} = 2-7$). Including the stellar component lowers $Q$ overall, but not enough to make $Q<1$ everywhere. The ring's global star formation efficiency (SFE) appears somewhat elevated relative to early spirals, but varies around the ring by more than an order of magnitude, peaking where star formation is most intense. AM0644-741's star formation law is peculiar: HI follows a Schmidt law while H$_2$ is uncorrelated with star formation rate density. Photodissociation models yield low volume densities in the ring, particularly in the starburst quadrant  ($n \\approx 2$ cm$^{-3}$), implying a warm neutral medium dominated ISM.  At the same time, the ring's pressure and ambient far-ultraviolet radiation field lead to the expectation of a predominantly molecular ISM. We argue that the ring's peculiar star formation law, $n$, SFE, and $f_{\\rm mol}$ result from the ISM's $\\ga100$ Myr confinement time in the starburst ring, which amplifies the destructive effects of embedded massive stars and supernovae. As a result, the ring's molecular ISM becomes dominated by small clouds where star formation is most intense, causing $M_{\\rm H_2}$ to be underestimated by $^{12}$CO line fluxes: in effect $X_{\\rm CO} \\gg X_{\\rm Gal}$ despite the ring's solar metallicity. The observed large HI component is primarily a low density photodissociation product, i.e., a tracer rather than a precursor of massive star formation. Such an ``over-cooked'' ISM may be a general characteristic of evolved starburst ring  galaxies. ",
        "introduction": "A large and growing body of evidence supports the key role played by gravitational interactions in the formation and evolution of galaxies. In the hierarchical structure formation framework, galaxies are assembled over cosmic time through multiple mergers \\citep[e.g.,][]{white,bell06,jogee}. Evidence for this can be found in the observed asymmetric, sometimes chaotic morphologies of a significant fraction of galaxies detected in deep surveys \\citep[e.g.,][]{williams,ferguson,conselice}. Similarly, there are indications that submillimeter galaxies, thought to be the precursors of current epoch massive bulge dominated galaxies, experience star formation rates (SFR) in excess of $\\approx1000$ $M_{\\odot}$ yr$^{-1}$ as a direct result of mergers \\citep[e.g.,][]{smail,richards,tacconi}. Morphological evolution is also clearly attributable to close passages and mergers \\citep[e.g.,][]{toomre74,dressler,moore,hibbard}. Interactions likewise appear involved with the funneling of material to the immediate vicinity of supermassive black holes and thus the switching on of active galactic nuclei \\citep[AGNs; e.g.,][]{heckman,ramos}. Detailed numerical and observational studies of strongly interacting systems in the local universe are thus highly relevant, especially when the dynamical history of the encounter can be reasonably well reconstructed, allowing detailed numerical simulations capable of following the evolution of star-forming interstellar medium (ISM) that can be constrained by observations.  Ring galaxies are striking examples of the ability of collisions to transform both the morphology and star-forming activity of a spiral galaxy. Numerical studies since the mid-1970s argue persuasively that the optically prominent rings are in fact outwardly propagating zones of strong orbit crowding within the disk of a spiral induced by the near central passage of a companion galaxy \\citep[e.g.,][but see Korchagin et al. 1999, 2001 for a very different formation mechanism]{lynds,struckhigdon,mapelli}. The expanding rings are transient but maintain their coherence for roughly one outer-disk rotation period ($\\sim$300 Myr). H$\\alpha$ emission line imaging show the rings to be nearly always sites of intense massive star formation, while star formation within the enclosed disks is greatly quenched \\citep[e.g.,][]{wof1,marston}.\\footnotemark[1] Ring galaxies are thus of interest from the standpoint of both the triggering and suppression of star formation on galactic scales. HI interferometry of the Cartwheel ring galaxy provides insight into how star formation is regulated in these objects \\citep{wof2}. Its outer ring possesses $\\approx90\\%$ of the galaxy's total neutral atomic ISM, resulting in very high HI surface densities ($\\Sigma_{\\rm HI}$), peaking at $\\Sigma_{\\rm HI} = 75$ $M_{\\odot}$ pc$^{-2}$ for 3 kpc linear resolution.\\footnotemark[2] At the same time, the Cartwheel's interior disk, inner ring, and nucleus are HI poor, with $\\Sigma_{\\rm HI} \\lesssim 2$ $M_{\\odot}$ pc$^{-2}$ ($3 \\sigma$). Star formation triggering/suppression can be understood in terms of a local gravitational stability criterion, as parameterized by the surface density of the neutral atomic ISM. The Cartwheel's outer HI ring exceeds the gas surface density threshold for robust star formation \\citep{kennicutt89}, while the enclosed gas poor disk is stable against the growth of large scale gravitational perturbations and subsequent star formation. Whether simple gravitational stability ideas can explain the distribution of star formation in other large and evolved ring galaxies remains to be seen. \\footnotetext[1]{One notable exception is the ``Sacred Mushroom'' galaxy (AM1724-622), whose red and low surface brightness optical ring shows no evidence of recent star formation, and indeed suggests that the pre-collision galaxy was an S0 \\citep{wallin}.} \\footnotetext[2]{HI surface density and column density are related: 1 $M_{\\odot}$ pc$^{-2}$ = 1.25 $\\times$ 10$^{20}$ atoms cm$^{-2}$.}  Stars form in the cold molecular ISM, so HI only tells part of the story. A more complete understanding of star formation triggering and suppression in ring galaxies requires the determination of the {\\em molecular} ISM's properties as well. This has proven difficult since ring galaxies as a rule are not luminous in the $^{12}$CO rotational transitions \\citep[e.g.,][]{hor95}, which are the standard tracers of H$_2$ in galaxies. This has been attributed to low metal abundances in the ring from the ``snow-plowing'' of largely unenriched gas from the outer disk as the ring propagates outward. In this paper we examine the atomic and molecular ISM of the southern ring galaxy AM0644-741 \\citep{arp}, also known as the  ``Lindsay-Shapley Ring'', using data from the Australia Telescope Compact Array (ATCA)\\footnotemark[3] and the 15 m Swedish ESO Submillimeter Telescope (SEST)\\footnotemark[4] at La Silla, Chile. $B$-band and H$\\alpha$ images are shown in Figure 1 and a summary of AM0644-741's global properties is given in Table 1. The ring is very large ($D_{\\rm ring} = 95\\arcsec$, or $42$ kpc at the assumed distance), with only those of the Cartwheel and AM1724-622 (the ``Sacred Mushroom'') being larger \\citep[][]{wof1,wof2,wallin}. Optical spectroscopy by \\citet[][hereafter FMA]{fma} showed the ring to be rich in giant star-forming complexes, and expanding in accord with collisional models. H$\\alpha$ imaging by \\citet[][hereafter HW97]{wof3} showed massive star formation to be confined to a pair of (apparently) interlocking rings, with a total SFR$=3.6$ $M_{\\odot}$ yr$^{-1}$ with no extinction correction. (An unresolved nuclear source contributes $\\approx1\\%$ of the total $F_{\\rm H\\alpha}$.) The offset nucleus and asymmetric distribution of H$\\alpha$ emission point to an off-centered collision with the elliptical $1.8\\arcmin$ ($49$ kpc in projection) to the southeast in Figure 1. $N$-body/smoothed particle hydrodynamics (SPH) models of this system have been presented in \\citet{antunes}. \\citet{hor95} reported the detection of $^{12}$CO($J=1-0$) emission from the galaxy's nucleus using the SEST (though at a velocity $300$ km s$^{-1}$ smaller than $v_{\\rm sys}$). While nuclear $^{12}$CO emission does not guarantee the detection of molecular gas in the starburst rings \\citep[e.g.,][]{hor98}, it motivated the work presented in this paper. \\footnotetext[3]{The Australia Telescope Compact Array is part of the Australia Telescope, which is funded by the Commonwealth of Australia for operations as a National Facility managed by CSIRO.} \\footnotetext[4]{The SEST is operated jointly by the European Southern Observatory (ESO) and the Swedish National Facility for Radio Astronomy, Chalmers Center for Astrophysics and Space Science.}  We use the ATCA HI data to trace the distribution and kinematics of AM0644-741's neutral atomic ISM on $\\approx 3$ kpc scales, which is sufficient to allow direct comparisons with published VLA HI observations of the Cartwheel. The radio continuum data additionally provides an obscuration free tracer of current massive star formation. The SEST is used to measure both 115 GHz $^{12}$CO(J = 1-0) and 230 GHz $^{12}$CO(J = 2-1) line emission in the ring galaxy. We take advantage of improved receiver sensitivity, particularly at $230$ GHz, to reexamine the galaxy's molecular gas content. A major advantage of 230 GHz observations is the $22\\arcsec$ FWHM beam, which allows the measurement of $^{12}$CO emission in the ring, disk, and nucleus separately, permitting the best census of the molecular ISM in a robustly star-forming ring galaxy in the pre-ALMA era.  We also utilize far-ultraviolet (FUV) images of AM0644-741 obtained with the Galaxy Evolution Explorer (GALEX), ground-based optical imaging obtained at the Cerro-Tololo Inter-American Observatory (CTIO)\\footnotemark[5], and images at 4.5, 24, and 70 $\\mu$m from the Spitzer Space Telescope \\citep[$Spitzer$,][]{werner}\\footnotemark[6], as well as optical spectra of AM0644-741's ring using ESO's 3.6 m telescope. The combined data are used to determine basic properties of the neutral atomic and molecular ISM throughout AM0644-741's ring, nucleus and enclosed disk for comparison with star formation activity as measured in the FUV, optical, infrared, and radio. Among the questions we wish to address is whether the ring's neutral ISM is dominated by atomic or molecular gas, and which local property is primarily responsible for setting the molecular fraction. We also wish to know if the distribution of star formation in AM0644-741's ring and disk can be understood in terms of simple gravitational stability arguments as is the case in the Cartwheel, and whether or not the inclusion of the molecular ISM significantly modifies the results of such an analysis. H$_2$ mass measurements are also necessary for estimates of star formation efficiency (SFE), which have been determined in nearby galaxies \\citep[e.g.,][]{young,rownd}. In short, how do properties of the star-forming ISM in an evolved ring galaxy compare with other systems? \\footnotetext[5]{CTIO is operated by AURA, Inc., under contract with the National Science Foundation.} \\footnotetext[6]{Spitzer Space Telescope is operated by JPL, California Institute of Technology for the National Aeronautics and Space Administration.}  This paper is organized as follows: The observations and data reduction are outlined in Section 2. The results derived using these observations are presented in Section 3, including the ring's metallicity, the distribution and kinematics of HI, the relative proportions of molecular and atomic gas, star formation efficiency, and their variations around the ring. We also consider the ring's gravitational stability, and whether or not star formation follows a Schmidt law. In Section 4 we consider the role played by pressure and ambient FUV radiation field in determining the ring's molecular fraction. We also discuss our results in the context of an alternate view of ring galaxies, in which the ring represents a self-sustaining propagating starburst. Our conclusions are given in Section 5. Throughout this paper we adopt a flat $\\Lambda$CDM cosmology based on the $7$ year WMAP results \\citep{wmap7}, $\\Omega_{\\rm M}$=0.27, $\\Omega_{\\rm \\Lambda}$=0.73, and a Hubble constant of 70.3 km s$^{-1}$ Mpc$^{-1}$, giving AM0644-741 a luminosity distance of $96.9$ Mpc and a linear scale of $0.45$ kpc arcsec$^{-1}$. ",
        "conclusions": "Ring galaxies are unique systems in which to study the triggering and regulation of massive star formation on large scales. In this paper we examine one of the best examples, AM0644-741, which is notable for its large diameter ring ($42$ kpc), inferred age ($133$ Myr), and global SFR. From the analysis of H$\\alpha$, FUV, infrared, and radio continuum images we conclude that SFR$ = 11$ $M_{\\odot}$ yr$^{-1}$, with relatively low internal extinction ($A_{\\rm H\\alpha}\\approx1$), and no significant star formation within the encircled disk or nucleus. Analysis of optical spectra show that AM0644-741's ring possesses $\\ga$solar metal abundances (in marked contrast with its ``twin'' the Cartwheel), opening the possibility of characterizing its molecular ISM using rotational transitions of $^{12}$CO. The distribution of HII complexes gives the appearance of the classic ``beads on a string'' morphology, which is a hallmark of star formation triggered by the action of large scale gravitational instabilities. One therefore expects $Q<1$ throughout the ring.  Observations using ATCA and SEST show that essentially all of AM0644-741's neutral ISM is concentrated in the expanding $42$ kpc diameter ring, which agrees both with collisional models and high resolution observations of other ring galaxies (e.g., Arp~143 and the Cartwheel). Contrary to our expectations, we find the ring's global molecular fraction to be surprisingly small ($f_{\\rm mol} = 0.078 \\pm 0.003$), particularly given that Sa-Sb galaxies, the likely progenitor of AM0644-741, normally possess comparable amounts of HI and H$_2$. It is this apparent deficiency of H$_2$ that is responsible for the ring's other peculiar properties: elevated SFE, a star formation law where atomic gas - but not molecular gas - correlates with star formation, and typically $Q_{\\rm tot}\\ga1$. The concentration of the disk's ISM into the dense ring appears to have not generated a molecular rich environment. Quite to the contrary, applying PDR models to the observed $\\Sigma_{\\rm HI}$, $\\Sigma_{\\rm H_2}$, and $\\chi_{\\rm UV}$ yield WNM-like total hydrogen volume densities ($n\\approx2$ cm$^{-3}$) throughout the ring quadrant experiencing peak SFR.  Our data allow us to reject the possibility that the ring's apparently small molecular component is the result of low and/or variable metallicity, unusually high $\\chi_{\\rm UV}$, or insufficient pressure (which regulates the conversion of HI to H$_2$). We find the opposite to be true: AM0644-741's ring would have been expected to be dominated by molecular gas given our estimates of its $P_{\\rm ISM}$ and $\\chi_{\\rm UV}$. To explain how AM0644-741 can form stars in a WNM dominated ISM, we propose that the rings in evolved and robustly star-forming ring galaxies like AM0644-741 and the Cartwheel possess an ``over-cooked'' ISM, where GMCs are fragmented and photodissociated by prolonged exposure to SNe shocks, expanding HII regions, and moderately high $\\chi_{\\rm UV}$ ($\\approx 7-25 \\chi_{\\odot}$) during their $>100$ Myr confinement to the ring. Consequently, CO survives only in the small innermost cloud cores due to reduced shielding available in the cloud fragments, and traces H$_2$ differently than for a Galactic cloud. In particular, $^{12}$CO transitions will significantly {\\em underestimate} the amount of H$_2$ in the ring in a way that depends on the local SFR density. An ``over-cooked'' ISM can thus account for the unusually high SFE, the small $f_{\\rm mol}$ and its variations around the ring, as well as the peculiar star formation law, where CO appears to be decoupled from local star formation. This also explains the observed linear relationship between $\\Sigma_{\\rm SFR}$ and $\\Sigma_{\\rm HI}$: the abundant atomic gas we detect with the ATCA is primarily a photodissociation product, which traces rather than fuels star formation. It is no longer surprising, then, that this component will resemble the WNM. For all practical purposes, it is the WNM.  If, as we suspect, $^{12}$CO significantly underestimates the ring's molecular component, then $Q_{\\rm tot}$ may be significantly overestimated. If the H$_2$ obeys an M51-like Schmidt Law (e.g., K07), then we find that $Q_{\\rm tot}<1$ is satisfied for nearly all of the ring, and in particular, where massive star formation is observed to be robust. Gravitational instabilities then become the dominant star formation trigger. Independent estimates of the ring's molecular component, as well as better constraints on the stellar velocity dispersion, are clearly needed to conclusively settle this point.  Another peculiar aspect of AM0644-741's ISM, and one not evident in the Cartwheel, is the wide and multi-component line profiles, both in the atomic and molecular gas. It is not clear at present if this reflects caustics in the gas flow or out-of-plane motions, though preliminary $N$-body simulations appear to favor the former interpretation (J. Wallin, private communication).  Our results imply that the rings of systems like AM0644-741 and the Cartwheel represent substantially different star-forming environments than that typically found in spiral arms due to the sustained accumulation of damage to the ISM from concentrated populations of OB stars. At the same time, the modest FUV radiation fields mean that they are still different from intense nuclear starbursts and ULIRGs. Ring galaxies remain interesting and valuable systems to study the interaction between the ISM and intense star formation, and follow-up observations with $Herschel$ and ALMA, in addition to models incorporating more sophisticated treatments of the ISM, are clearly warranted.  Finally, we examined the ``forest fire'' interpretation of ring galaxies in light of these (and previous) observational results. We argue that a propagating starburst \\citep[e.g.,][]{gerola} cannot reproduce the observed gas distribution or large-scale kinematics of ring galaxies. Likewise, the space density of nuclear starburst galaxies, i.e., systems possessing the forest fire ``detonator'', is $\\approx50$ times larger than that of ring galaxies, implying that if ring galaxies arise this way, they should be much more common. The collisional interpretation has no such problems."
    },
    "1107/1107.2166_arXiv.txt": {
        "abstract": "We analyze the a set of seventeen rotation curves of Low Surface Brightness (LSB) galaxies from the The HI Nearby Galaxy Survey (THINGS) with different mass models to study the core structure and to determine a phase transition energy scale ($E_c$) between hot and cold dark matter,  due to nonperturbative  effects in the Bound Dark Matter (BDM) model.   Our results agree with previous ones implying the cored profiles are preferred over the N-body motivated cuspy Navarro-Frenk-White (NFW) profile.  We find an average galactic core radius of $r_c = 260\\times 10^{\\pm 1.3}$ pc and a phase transition energy $E_c = 0.11\\times 10^{\\pm  0.46}{\\rm \\ eV}$, that is of the same order of magnitude as the sum of the neutrino masses. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.4095_arXiv.txt": {
        "abstract": "We simulate mergers between galaxies containing collisionally-relaxed nuclei around massive black holes (MBHs). Our galaxies contain four mass groups, representative of old stellar populations; a primary goal is to understand the distribution of stellar-mass black holes (BHs) after the merger. Mergers are followed using direct-summation $N$-body simulations, assuming a mass ratio of 1:3 and two different orbits. Evolution of the binary MBH is followed until its separation has shrunk by a factor of 20 below the hard-binary separation. During the galaxy merger, large cores are carved out in the stellar distribution, with radii several times the influence radius of the massive binary. Much of the pre-existing mass segregation is erased during this phase. We follow the evolution of the merged galaxies for approximately three, central relaxation times after coalescence of the massive binary; both standard, and top-heavy, mass functions are considered. The cores that were formed in the stellar distribution persist, and the distribution of the stellar-mass black holes evolves against this essentially fixed background. Even after one central relaxation time, these models look very different from the relaxed, multi-mass models that are often assumed to describe the distribution of stars and stellar remnants near a massive BH.  While the stellar BHs do form a cusp on roughly a relaxation time-scale, the BH density can be much smaller than in those models. We discuss the implications of our results for the EMRI problem and for the existence of Bahcall-Wolf cusps. ",
        "introduction": "\\label{sec:intro}  The massive black holes (MBHs) that reside at the centers of some nearby galaxies are believed to grow together with their hosts through mergers: MBHs grow partly as a result of gas accretion, and partly by coalescence with other MBHs that are brought into the nucleus during the merger process \\citep{BBR80}.  The detailed assembly history of MBHs is poorly understood; major uncertainties include the ``seed'' mass distribution of MBHs at high redshift, the typical gas accretion efficiency, and the frequency with which MBHs are ejected due to gravitational-wave recoil \\citep{KH00,VHM03}.  But a robust prediction of the hierarchical models is that galaxies hosting MBHs in the nearby universe were formed from less massive systems, at least some of which already contained MBHs.  Binary MBHs created during galaxy mergers leave imprints on the stellar distribution: for instance, they create low-density cores, by exchanging energy with passing stars \\citep{BBR80,mm01}.  Such cores are observed to be ubiquitous in stellar spheroids brighter than $\\sim 10^{10}L_\\odot$ \\citep{ferr1994,Lauer95} and their sizes -- of order the influence radius of the (presumably single) MBH -- are consistent with the predictions of merger models \\citep{graham2004,merritt06}. Here we define the influence radius as \\begin{equation} \\label{eq:rh} \\rh \\equiv \\frac{G\\,M_\\bullet}{\\sigma^2} \\end{equation} where $\\sigma$ is the rms velocity of stars in any direction at $r\\gap \\rh$.  Cores of radius $\\sim \\rh$ become difficult to resolve in galaxies beyond the Local Group if the MBH mass is below $\\sim 10^8\\msun$.  Even in the nearest nucleus, that of the Milky Way, the presence of a parsec-scale core around Sgr A$^*$ was only clearly established in the last few years \\citep{buchholz09}.  In the absence of MBHs, mergers tend to preserve the form of the stellar distribution near the centers of galaxies \\citep{Dehnen2005}. Binary MBHs, however, are efficient at erasing the structure that was present on scales $\\lap\\rh$ \\citep{mc2001}, and this fact precludes drawing definite conclusions about the nuclear properties of the galaxies that preceded the hosts of observed MBHs.  On the other hand, it is well established that low-luminosity galaxies have higher central densities than high-luminosity galaxies \\citep{Kormendy1985}. This is true both in terms of the mean density within the half-mass radius, and also in terms of the density on the smallest resolvable scales: low-luminosity spheroids often contain dense, nuclear stars clusters (NSCs), with sizes of order $10\\pc$ and masses $\\lap 1\\%$ the mass of the galaxy \\citep{Boeker2010}.  NSC masses are therefore comparable to, or somewhat greater than, MBH masses \\citep{Ferrarese2006}, although NSCs have been shown to coexist with MBHs in only a handful of galaxies \\citep{Seth2008,GS2009}.  The NSC in the Milky Way is believed to be a representative example: its half-light radius is $3-5\\pc$, or $1-2\\rh$, and its mass is a few times $10^7\\msun$, or several times $M_\\bullet$ \\citep{GS2009,Schoedel2011}.  In the NSC of the Milky Way, the two-body relaxation time is comparable with the age of the universe, and this is consistent with the persistence of a core (as opposed to a \\citet{BW76} cusp) in the late-type stars \\citep{merritt2010}.  But in fainter systems, central relaxation times are shorter.  For instance, in the Virgo cluster, galaxies with NSCs have nuclear half-light relaxation times that scale with host-galaxy luminosity as \\citep{merritt2009} \\beq T_{r,\\mathrm{NSC}}\\approx 1.2\\times 10^{10} \\mathrm{yr} \\left(\\frac{L_\\mathrm{gal}}{10^{10}L_\\odot}\\right).  \\eeq By comparison, the mean time between ``major mergers'' (mergers with mass ratios $3:1$ or less) of dark-matter haloes in the hierarchical models varies from $\\sim 0.2$ Gyr at redshift $z=10$ to $\\sim 10^{10}$ yr at $z=1$, with a weak dependence on halo mass \\citep{Fakhouri2010}.  This comparison suggests that the progenitors of many spheroids in the current universe may have been galaxies containing nuclei that were able to attain a collisionally-relaxed state before the merger that formed them took place.  In the absence of a MBH, collisional relaxation implies mass segregation and core collapse.  If a MBH is present, mass segregation still occurs, but core collapse is inhibited by the fixed potential due to the MBH.  A collisional steady state, which is reached by a time $\\sim T_r(\\rh)$ at radii $r\\lap \\rh$, is characterized by a Bahcall-Wolf, $n\\sim r^{-7/4}$ cusp in the dominant component at $r\\lap 0.2\\rh$.  If there is a mass spectrum, less-massive objects follow a shallower profile, $n\\sim r^{-3/2}$, while more-massive objects follow a steeper profile, $n\\sim r^{-2}$ \\citep{BW77,HA06}.  As a first approximation, the mass spectrum of an evolved stellar population can be represented in terms of just two components: objects of roughly one solar mass or less (main-sequence stars, white dwarfs, neutron stars); and remnant black holes (BHs) with masses $10-20\\msun$.  Standard initial mass functions predict that roughly $1\\%$ of the total mass will be in stellar BHs \\citep{Alexander2005}; so-called ``top-heavy'' mass functions \\citep[e.g.][]{Bartko10} predict a larger fraction.  In a collisionally relaxed nucleus, the density of stellar BHs will rise more steeply toward the center than the density of the stars.  Mergers between galaxies with such nuclei would be expected to modify these steady-state distributions substantially, and also to affect (increase) the time scale over which a collisionally relaxed cusp could be regenerated following the merger \\citep{MS06}.  These arguments motivated us to carry out merger simulations between galaxies containing multi-component, mass-segregated nuclei around MBHs.  As in previous papers from this series \\citep{mm01,ber05,mms07}, our merger simulations are purely stellar-dynamical.  In some galaxies, torques from gas would assist in the evolution of binary MBHs \\citep{escala05,dotti07,cuadra09}.  Gas also implies star formation, and there is evidence for complex star formation histories in many NSCs \\citep{Walcher2006}.  But $N$-body simulations that allow for non-spherical geometries \\citep{ber06,khan2011} have shown that purely dissipationless energy exchange with ambient stars can bring binary MBHs to milliparsec separations on time scales much shorter than galaxy lifetimes. Unless the late evolution of the binary MBH is greatly accelerated by torques from the gas, the influence of the binary on the distribution of the {\\it stellar} populations should be accurately reproduced by our dissipationless models.  With respect to star formation, population synthesis of NSC spectra suggest that most of the mass typically resides in stars with ages of order 5-10 Gyr \\citep{figer2004,Boeker2010}, i.e. an old population.  Simulating the merger of galaxy-sized systems, while enforcing the spatial and temporal resolution required to faithfully reproduce the dynamics of stars on scales $\\ll\\rh$ around the central MBH, is computationally demanding.  Our simulations used $\\sim 10^6$ particles per galaxy, and the models were advanced using a parallel, direct-summation $N$-body code \\citep{Harfst2007}.  The integrations were accelerated using special-purpose hardware.  The galaxy models contained four mass groups, representing an evolved stellar population.  Initial conditions of the merging galaxies were constructed in a two-stage process: models of mass-segregated nuclei around a MBH were first constructed, then these collisionally-relaxed models were imbedded into larger, spheroid-sized models.  Mergers were then carried out, assuming a galaxy mass ratio of $1:3$.  A major motivation for our new simulations was the need to understand the distribution of stellar remnants, particularly stellar-mass BHs, near the centers of galaxies.  Knowledge of the BH density well inside $\\rh$ is crucial for predicting the rates of many astrophysically interesting processes; in particular, the rate of capture of stellar-mass BHs by MBHs, or EMRIs \\citep{EMRIreview}.  Published EMRI rate calculations almost always assume a state of mass segregation, implying a high density of stellar remnants near the MBH \\citep{HA06,Hopman2009}. However, in a nucleus formed via a merger, any pre-existing mass segregation would have been disrupted by the binary MBH when it created a core; whether or not the massive remnants would have had time to re-segregate following the merger is difficult to assess without full $N$-body simulations.  Our simulations followed the evolution of the binary MBHs for a time almost long enough that gravitational wave emission would dominate the binaries' evolution.  We then combined the two MBH particles into one, simulating gravitational wave coalescence, and continued the $N$-body integrations for a time corresponding to several relaxation times at the (new) influence radius.  In this post-merger evolutionary phase, we also considered the consequences of varying the relative numbers of the different mass components.  In this way, we were able, for the first time, to observe how rapidly the stellar BHs would re-segregate following a merger.  We found substantially longer time scales for this evolution than in earlier simulations that started from physically less motivated initial conditions.  The paper is organized as follows. In Section~\\ref{sec:models} we describe the procedure to generate equilibrium segregated models starting from single-component models. Scaling to physical units is discussed in Section~\\ref{sec:scaling}. Evolution of the binary MBH and its effects on the underlying stellar distribution are described in Section~\\ref{sec:premerger}. In Section~\\ref{sec:postmerger} we describe the evolution of the light and heavy objects after the massive binary has undergone coalescence.  Section~\\ref{sec:shape} describes the shapes and kinematics of the merger remnants. Section~\\ref{sec:disc} discusses the implications of our results for the formation and observability of Bahcall-Wolf cusps, and for the distribution of stellar remnants. ",
        "conclusions": "\\label{sec:disc}  \\subsection{(Re-)growth of Bahcall-Wolf cusps in multi-component systems} In galaxies containing a single stellar population, a $\\rho\\propto r^{-7/4}$ \\citet{BW76} cusp is expected to appear at radii $r\\le r_\\mathrm{BW}\\approx 0.2 \\ri$ around the MBH.  While the growth time of the cusp is dependent on the initial conditions, simulations starting from a shallow cusp inside $\\ri$ show that the stellar density will have reached an approximate steady state after roughly one relaxation time at $\\ri$. We presented an example of such evolution in Figure\\,\\ref{fig:C2}.  In nuclei with two mass groups, e.g. solar-mass stars (MS) and $10\\msun$ black holes (BHs), evolution toward a steady state near the MBH depends on the relative numbers and masses in the two groups, as well as on the initial conditions.  The results of our $N$-body and Fokker-Planck integrations of two-component models were presented in Section~\\ref{sec:FP}.  Figure\\,\\ref{fig:FP_g0.5_slopes} showed that addition of the BHs reduces slightly the rate of formation of a Bahcall-Wolf cusp in the MS (light) component, and also affects the final value of the density-profile slope, which varies from $\\rho_\\mathrm{MS}\\sim r^{-7/4}$ when $\\rho_\\mathrm{BH}/\\rho_\\mathrm{MS}$ is small, to $\\sim r^{-3/2}$ as $\\rho_\\mathrm{BH}/\\rho_\\mathrm{MS}$ is increased \\citep{BW77}.  In addition, if the initial MS distribution is very flat near the MBH, scattering by the heavier BHs can dominate the evolution of the MS component at early times for $r\\lap 0.05 \\ri$, converting $\\rho_\\mathrm{MS}\\sim r^{-1/2}$ to $\\rho_\\mathrm{MS}\\sim r^{-3/2}$ at these small radii, even before the Bahcall-Wolf cusp has fully formed at larger radii.  The full galaxy merger simulations presented here allowed us, for the first time, to evaluate the evolution of multi-component nuclei starting from initial conditions that were motivated by a well-defined physical model.  Our pre-merger galaxies contained mass-segregated nuclei with four mass components, representing an evolved stellar population.  These initial distributions were modified both by the galaxy merger, and by the formation of a MBH binary, which created a large core in each of the components (Figures\\,\\ref{fig:lagr1}, \\ref{fig:density}).  The core radius of the heaviest (BH) component was somewhat smaller than the cores in the three lighter components, a relic of the earlier mass segregation, and of the incomplete cusp destruction process during the binary MBH phase (Figure\\,\\ref{fig:density}).  Evolution during the merger phase set the ``initial conditions'' of the nucleus at the time when the two MBHs were combined into one.  Unlike the rather ad hoc initial conditions used in many earlier studies of nuclear evolution \\cite[e.g.][]{Freitag06,PAS2010}, our post-merger models have cores that should be reasonable representations of the cores in real galaxies that formed via dissipationless mergers, with mass densities that decline toward the center and anisotropic kinematics that reflect the action of the massive binary on the stellar orbits (Figure\\,\\ref{fig:vel}).  By continuing the evolution of these multi-component models for a time greater than $T_r(\\ri)$ after coalescence of the two MBHs, we found that the cores characterizing the distribution of the dominant, MS component persisted; in fact the MS density at radii $\\sim \\ri$ decreased gradually with time (Figure~\\ref{fig:massp}) -- a predictable consequence of the initial extent of the cores, several times $\\ri$, implying $T_r(r_c)> T_r(\\ri)$, and of continued ``heating'' by the heavier BHs.  Nevertheless, a Bahcall-Wolf cusp gradually reformed in the lighter components at radii $r\\lap 0.2\\ri\\ll r_c$; growth times were found to be  $\\gap T_r(\\ri)$, somewhat greater than in earlier models with more idealized initial conditions \\cite[e.g.][]{Freitag06}.  Our pre-merger galaxies (models A1, B1) contained larger numbers of remnants than would be expected based on a standard IMF; this was done in order to better resolve the distributions of those components near the MBH.  We tested the dependence of the {\\it post}-merger evolution on the assumed mass function by carrying out a second set of integrations in which we decreased the relative numbers of remnants (NS, WD, BH) by factors of a few, to values more consistent with standard IMFs (models A2, B2).  Evolution of the dominant, MS component in the latter models differed only modestly from its evolution in the models with larger remnant fractions; the main difference was a lower rate of core expansion reflecting a lower rate of heating by the BHs (Figure\\,\\ref{fig:massp}).  The rate of growth of the Bahcall-Wolf cusp in the MS component was essentially unchanged (Figure\\,\\ref{fig:profiles}).  Our results on the regeneration of Bahcall-Wolf cusps following dissipationless mergers are consistent with those obtained in simulations of single-component galaxies \\citep{MS06}.  We can summarize these results by stating that regrowth of a cusp in the dominant stellar component requires a time comparable with, or somewhat longer than, the relaxation time of that component measured at the MBH influence radius.  The new simulations presented here suggest that time scales for cusp regrowth are only weakly dependent on the number of heavy remnants (BHs), at least if the fraction of mass in the heavier population does not exceed a few percent of the total.  The Milky Way nucleus is near enough that a Bahcall-Wolf cusp could be resolved if present, and the dominant stellar population is believed to be old.  Since $\\rh\\approx \\ri\\approx 2-3\\pc$ in the Milky Way, the expected, outer radius of the Bahcall-Wolf cusp is $r_\\mathrm{BW}\\approx 0.5\\pc \\approx 10''$.  As is well known, number counts of the dominant, old stellar population show no evidence of a rise in density at this radius; instead the number counts are flat, or even falling, from $\\sim 10''$ into at least $1''$ projected radius \\citep{buchholz09,Do09,Bartko10}.  Assuming solar-mass stars, the relaxation time at the influence radius of Sgr A$^*$ is 20-30 Gyr \\citep{merritt2010}, while the mean stellar age in the nuclear star cluster is estimated to be $\\sim 5$ Gyr \\citep{figer2004}.  The time available for formation of a cusp is therefore $(5/25) T_r(\\ri)\\approx 0.2T_r(\\ri)$.  Our simulations (e.g. Figure\\,\\ref{fig:profiles}) suggest that a Bahcall-Wolf cusp in the dominant component is unlikely to have formed in so short a time, and this is consistent with the lack of a Bahcall-Wolf cusp at the Galactic center.  The core in the Milky Way has a radius of $\\sim 0.5\\pc$, somewhat smaller than $\\rh$ or $\\ri$, while the cores formed in our merger models are somewhat larger than $\\ri$ (Table 4).  It has been argued that the Milky Way core is small enough that gravitational encounters would cause it to shrink appreciably in 10 Gyr, as the stellar distribution evolves toward a Bahcall-Wolf cusp \\citep{merritt2010}.  The cores in our $N$-body models are so large that they do not evolve appreciably after the binary MBH has been replaced by a single MBH; the Bahcall-Wolf cusp forms at radii smaller than $r_c$, leaving the core structure essentially unchanged.  Without necessarily advocating a merger model for the origin of the Milky Way core (it is unclear whether our galaxy even contains a bulge; \\citet{Valpuesta2011}), we note that the sizes of cores formed by binary MBHs scale with the binary mass ratio. Presumably, we could have produced cores more similar in size to the Milky Way's if we had adjusted this ratio.  It has been suggested \\citep{LBK2010} that heating by BHs could be responsible for the lack of a Bahcall-Wolf cusp at the Galactic center.  We do not find support for this hypothesis, either in our four-component $N$-body models, nor in our two-component Fokker-Planck models.  The latter showed that even a quite large BH population still allowed a Bahcall-Wolf cusp to form in the MS stars on a time scale of $\\sim T_r(\\ri)$; the main effect of the BHs is to decrease the asymptotic slope of the cusp from $\\sim r^{-7/4}$ to $\\sim r^{-3/2}$.  \\subsection{The distribution of massive remnants in galaxy nuclei}  A dense cluster of stellar-mass BHs has been invoked as a potential solution to a number of problems of collisional dynamics at the Galactic center.  Examples include randomization of the orbits of young stars via gravitational scattering \\citep[e.g.][]{perets09}, production of hyper-velocity stars through encounters with BHs \\citep[e.g.][]{OL08}, and warping of young stellar disks \\citep{KT11}. These treatments typically assume a collisionally-relaxed state for the Galactic center.  In the relaxed models, the mass in BHs inside $0.1\\pc$ is $\\sim 10^4\\msun$, i.e. $N_\\mathrm{BH} (r<0.1\\pc) \\approx 10^3$ and $N_\\mathrm{BH}(r<0.01\\pc)\\approx 10^2$ \\citep{HA06,Freitag06}, assuming ``standard'' IMFs.  A high density of stellar BHs at the centers of galaxies like the Milky Way is also commonly assumed in discussions of the EMRI (extreme-mass-ratio inspiral) problem \\citep{EMRIreview}.  Models like these are called into question by the lack of a Bahcall-Wolf cusp in the late-type stars at the center of the Milky Way \\citep{buchholz09,Do09,Bartko10}.  If the Galactic center is not collisionally relaxed, as these observations seem to suggest, then computing the distribution of the heavy remnants becomes a more difficult, time-dependent problem \\citep{merritt2010}, and knowledge of the initial conditions is essential.  Dissipationless mergers imply ``initial conditions'' characterized by a core in the dominant stellar component.  The cores formed in our (major) merger simulations are large enough that they do not evolve appreciably (i.e. shrink) even after $\\sim 3$ relaxation times at the MBH influence radius.  As a result, evolution of the BH distribution takes place against a stellar background with a very different density profile than in the steady-state models.  In a core around a MBH, dynamical friction is much weaker than one would estimate by plugging the local density into standard formulae for orbital decay, due to the absence of stars moving more slowly than the local circular velocity \\citep{AntoniniMerritt2011}.  Scaling our $N$-body models to the Milky Way, we predicted numbers of BHs inside $\\rh$ that were substantially smaller than in the collisionally relaxed models; in the inner $10^{-2}\\pc$ -- the radii most relevant to the EMRI problem \\citep{MAMW2010} -- the number of BHs was, at most, 10-100 times smaller than predicted by these models, even after 10 Gyr (Figure\\,\\ref{fig:BHs}).  It is unclear how relevant merger models are to the center of the Milky Way.  If the core observed at the Galactic center had some other origin, the distribution of stellar BHs might have little connection with the distribution of the giant stars.  However, if cores of size $r_c\\approx\\rh$ are common features of galactic nuclei, and if at some early time both the BHs and the stars had a common core radius, our models imply that the distribution of BHs should be considered very uncertain, even in galaxies for which nuclear half-mass relaxation times are as short as the age of the universe."
    },
    "1107/1107.1812_arXiv.txt": {
        "abstract": "We report the discovery of a substellar companion to the nearby solar-type star HD~46588 (F7V, 17.9~pc, $\\tau\\sim3$~Gyr). HD~46588~B was found through a survey for common proper motion companions to nearby stars using data from the Wide-field Infrared Survey Explorer and the Two-Micron All-Sky Survey.  It has an angular separation of $79.2\\arcsec$ from its primary, which corresponds to a projected physical separation of 1420~AU. We have measured a spectral type of L9 for this object based on near-infrared spectroscopy performed with TripleSpec at Palomar Observatory. We estimate a mass of $0.064^{+0.008}_{-0.019}$~$M_\\odot$ from a comparison of its luminosity to the values predicted by theoretical evolutionary models for the age of the primary. Because of its companionship to a well-studied star, HD~46588~B is one of the few known brown dwarfs at the L/T transition for which both age and distance estimates are available. Thus, it offers new constraints on the properties of brown dwarfs during this brief evolutionary phase. The discovery of HD~46588~B also illustrates the value of the Wide-field Infrared Survey Explorer for identifying brown dwarfs in the solar neighborhood via their proper motions. ",
        "introduction": "L and T dwarf companions to nearby stars are valuable for studying the structure and evolution of brown dwarfs as well as their origin. Because the ages and metallicities of these companions often can be determined via their primaries, they can be used to calibrate atmospheric and evolutionary models \\citep{kir05,cus08,hel08,bur10b}. Meanwhile, theories for the formation of brown dwarfs can be tested with measurements of the fraction of stars that harbor wide low-mass companions and the frequency of tight binaries among these companions \\citep{rei01,bat03,sta09}.  Multiple strategies have been employed in identifying cool companions to stars in the solar neighborhood. One approach is to define a sample of stars and search for companions to them \\citep{bec88,nak95,pot02,low05,bil06,met06,mug06,luh07,luh11}. Alternatively, one can begin with a sample of known brown dwarfs and check whether they are companions to stars \\citep{bur00,kir01,wil01,sch03,burn09,fah10,fah11,gol10,sch10,zha10,day11}. Late-type companions also can be uncovered through a search for pairs of comoving objects across a large area of sky \\citep{rad08}. Proper motion measurements are essential in all of these strategies, either for the initial identification of candidate companions or for confirming candidates that were selected through photometry. These proper motions can be measured through dedicated imaging or archival data from wide-field surveys like the Two Micron All-Sky Survey \\citep[2MASS,][]{skr06}, the Sloan Digital Sky Survey \\cite[SDSS,][]{yor00}, and the United Kingdom Infrared Telescope Infrared Deep Sky Survey \\citep[UKIDSS,][]{law07}. A new source of astrometry has recently become available via the Wide-field Infrared Survey Explorer \\citep[WISE,][]{wri10}. This mission has performed the first all-sky survey at infrared (IR) wavelengths since 2MASS, thus providing an opportunity to measure proper motions for 2MASS sources across the entire sky. Indeed, proper motions incorporating WISE data have already been exploited for searching for free-floating L and T dwarfs \\citep{sch11,giz11}.  We have begun a search for faint common proper motion companions to nearby stars using astrometry from 2MASS and WISE. In this paper, we report the discovery of a substellar companion to the solar-type star HD~46588. We describe the methods that uncovered this object as a possible companion (\\S~\\ref{sec:search}) and present spectroscopy that confirms its cool nature (\\S~\\ref{sec:spec}). We then examine the physical properties of this new companion (\\S~\\ref{sec:prop}) and discuss its relevance to studies of L and T dwarfs (\\S~\\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}  We have presented the discovery of a widely-separated L9 companion to the nearby F7V star HD~46588. Among previously known companions, Gl584~C most closely resembles HD~46588~B in terms of the spectral types of the system components \\citep[G1/G3/L8,][]{kir01}. Other late-L/early-T companions to stars include Gl337~CD \\citep[T0,][]{wil01,bur05}, HD~203030~B \\citep[L7.5,][]{met06}, HN~Peg~B \\citep[T2.5,][]{luh07}, and LHS~2397aB \\citep[L7,][]{dup09}.  Brown dwarfs near the L/T transition have served as laboratories for studying the complex atmospheric changes that occur between L and T dwarfs \\citep{bur10a}. Observations of binaries have played a central role in these studies. For instance, systems with spectral types of late L or early T appear to have a high binary fraction compared to other types \\citep{bur06b,liu06}. This has been explained by rapid evolution of brown dwarfs through the L/T transition, which causes single transition objects to be rare and binaries containing earlier and later types to dominate systems that have composite types near the L/T boundary \\citep{bur06b,bur07}. These binaries often appear overluminous relative to other brown dwarfs at the same color or spectral type. HD~46588~B is somewhat overluminous in $M_{W2}$ vs.\\ $K_s-W2$ but appears near the lower envelope of the L/T dwarf sequence in $M_{W2}$ vs.\\ $W1-W2$ (Figure~\\ref{fig:cmd}). Thus, it does not show clear evidence of binarity. Regardless of whether it is single or binary, since its distance is known, HD~46588~B offers an opportunity to better constrain the relations between spectral type and absolute magnitudes (e.g., $M_{\\rm bol}$, $M_K$) near the L/T transition, which influence the modeling of binary statistics and surface densities of L/T objects as well as the identification of overluminous systems \\citep{bur07}. In addition, because of its companionship to a star, HD~46588~B has the best available age estimate of any known object at a spectral type of L9. By combining its age and luminosity with evolutionary models, we have been able to estimate its temperature, and thus constrain the temperature of the L/T boundary. The available measurements of $M_{\\rm bol}$, $M_K$, and temperature for HD~46588~B are consistent with those derived previously for companions with ages of $>$1~Gyr \\citep{kir01,dup09}, but these relations could be more tightly constrained by improving the accuracy of near-IR photometry for HD~46588~B.  Data from the WISE mission have been recently used to identify new free-floating late-type members of the solar neighborhood based on their colors \\citep{mai11,bur11,cus11,kir11} and proper motions \\citep{sch11,giz11}. Our discovery of HD~46588~B illustrates the added value of WISE for uncovering cool companions to nearby stars via their proper motions. This object also helps to complete the census of L dwarfs within 20~pc from the Sun \\citep{cru07}."
    },
    "1107/1107.0605_arXiv.txt": {
        "abstract": "We present the multi-wavelength observations of a flux rope that was trying to erupt from NOAA AR 11045 and the associated M-class solar flare on 12 February 2010 using space and ground based observations from TRACE, STEREO, SOHO/MDI, {\\it Hinode}/XRT and BBSO.  While the flux rope was rising from the active region, an M1.1/2F class flare was triggered nearby one of its footpoints. We suggest that the flare triggering was due to the reconnection of a rising flux rope with the surrounding low-lying magnetic loops. The flux rope reached a projected height of $\\approx0.15 R_{\\odot}$ with a speed of $\\approx$90 km s$^{-1}$ while the soft X-ray flux enhanced gradually during its rise. The flux rope was suppressed by an overlying field and the filled plasma moved towards the negative polarity field to the west of its activation site. We find the first observational evidence of the initial suppression of a flux rope due to a remnant filament visible both at chromospheric and coronal temperatures that evolved couple of days before at the same location in the active region. SOHO/MDI magnetograms show the emergence of a bipole $\\approx$12 h prior to the flare initiation. The emerged negative polarity moved towards the flux rope activation site, and flare triggering near the photospheric polarity inversion line (PIL) took place. The motion of the negative polarity region towards PIL helped in the build-up of magnetic energy at the flare and flux rope activation site. This study provides a unique observational evidence of a rising flux rope that failed to erupt due to a remnant filament and overlying magnetic field, as well as associated triggering of an M-class flare. ",
        "introduction": "Magnetic dynamo action in the solar interior continuously generates magnetic flux tubes that rise through the photosphere, add new magnetic flux systems to the photosphere, and finally fan out in the corona \\cite{sch1999,asc2004}.  Such magnetic flux systems in the active regions are aligned with overlying arch filament systems, emerging in the form of $\\Omega$/U-loops, or may exhibit ``sea-serpent\"--like shapes (\\inlinecite{asc2004} and references therein). The newly emerging and growing magnetic fluxes may compel the topological changes in the overlying coronal magnetic fields, which may involve magnetic reconnection processes and their consequence in the form of eruptions at large and small spatial-temporal scales. MHD simulations of such magnetic flux tubes have been performed in the subphotospheric layers \\cite{fan2001,fan2003,fan2004} and also in the corona \\cite{shibata1990}. Emergence of magnetic field and its reconnection with the pre-existing field is not only responsible for the occurrence of large-scale eruptive phenomena, but it may also be the triggering mechanism for small-scale solar transients and explosive events (\\opencite{roussev2001a},\\citeyear{roussev2001b}, \\citeyear{roussev2001c}). Recently, the emergence of magnetic flux rope from the convection zone to the corona and its interaction with the pre-existing fields have been modelled and simulated by various authors ({\\it e.g.}, \\opencite{arch2004}, \\citeyear{arch2005}, \\citeyear{arch2007}).  The emergence of a dipole is not a sufficient condition to trigger a flare, because the emerging flux region may be too small or may have a non-favourable orientation \\cite{martin1984,asc2004}. Pre-flare brightenings and flares probably may not have direct relevance with the emerging flux tubes in many cases. However, they may be associated with the coupled energy release processes such as current driven plasma micro-instabilities, etc.\\cite{karpen1986}. It is most likely that the magnetic flux emergence does not trigger flares directly, but may instead increase the magnetic complexity locally by adding new magnetic flux and, helicity, or by inducing  motions of sunspots. When this magnetic complexity crosses the critical threshold, flares and associated eruptions may be triggered \\cite{asc2004,sri2010}. However, flare models that are directly driven by flux emergence have also been reported \\cite{hey1977,asc2004}. Moreover, the rising process of  the magnetic field through the convection zone and its emergence through the photosphere, chromosphere and corona are not well understood. The exact reason of the activation of helicity and twist in the pre-existing magnetic flux ropes are also not well explored both in the observations and theory.  The magnetic flux from the sub-photospheric levels may also be associated with the helicity and twist, which leads to more stored magnetic energy and instabilities. Recently, more observational signatures became evident that one of the origins of this twist is at a near sub-photospheric level from which the flux tube swirls above the surface \\cite{bonet2008,wede2009}. These phenomena are also described by recent theoretical models \\cite{simon1997,shel2011a,shel2011b,fedun2011}. The solar flares are mainly distinguished by two categories, {\\it i.e.}, confined and eruptive flares, and both of them usually take place in rather complex morphology of overlying magnetic fields in the active regions. The helicity and instability generated in such complex magnetic fields are the most likely cause of stored excess energy in the active regions that is later released in the form of solar flares and initiates the related plasma eruptions. The emergence of such twisted and unstable magnetic flux tubes, or their activation and interactions in the active regions, may also trigger the flare energy release and associated eruptions as studied recently by ,{\\it e.g.}, \\inlinecite{pankaj2010a}, \\inlinecite{pankaj2010b}, \\inlinecite{pankaj2010c}, \\opencite{sri2010}. On the modelling side, most of the numerical simulations have used magnetic flux emergence and sunspot rotation caused by photospheric plasma flows to increase the magnetic helicity leading to the final eruption ({\\it e.g.}, \\opencite{amari2000}, \\opencite{torok2003}, \\citeyear{torok2005}).   \\inlinecite{shen2011} have reported the dynamics of a filament that failed in the eruption several times, and finally succeeded to erupt with the triggering of a C-class flare process and associated CME on 8 February 2010 from active region (AR) NOAA 11045. Before that work, the observational evidence of flux ropes failed in eruption has been reported as well \\cite{ji2003,alex2006,liu2009}. \\inlinecite{boris2001} have studied the critical height and compared it with the stability criterion of the filaments. However, most of the previous observational studies of failed eruption could not reveal the exact mechanism associated with it. In this paper, we will study the dynamics of a flux rope in the AR 11045 on 12 February 2010 that moved up higher in the corona but still failed in the eruption. We also study the energy build-up and release processes of an M-class flare associated with the activation of this flux rope. We find the first observational signature of the rising flux rope that is failed due to an overlying remnant filament. In Section 2 we will present the multi-wavelength observations. We will discuss the magnetic configuration and scenario of the event in Section 3. Discussion and Conclusions will be presented in the last section. ",
        "conclusions": "In the present paper, we showed a rare multi-wavelength observation of a rising twisted flux rope on 12 February 2010 from NOAA AR 11045. The flux rope was activated in the vicinity of a remnant filament structure and moved up in the corona as visible in H$\\alpha$, EUV and soft X-ray observations. An M-class flare was triggered as the leading edge of the brightening moved away from the active region with a (projected) speed of $\\approx$90 km s$^{-1}$.  A two-ribbon flare progressed well with the ascent of the flux rope.  Most likely, the lower part of the flux rope structure reconnected with the surrounding low-lying loop system causing the flare. The flux rope reached to the maximum height and then it was suppressed by the overlying fields as the bending of the rope was observed in TRACE EUV observations. An analysis of associated MDI magnetograms show the emergence of a dipole $\\approx$12 h prior to the flare and flux rope activation. The continuous growth and motion of two spots in the opposite direction (particularly the negative polarity spot), helped in building the magnetic energy necessary to initiate a flare in the vicinity of the negative spot, where the flux rope activation took place. The structure moved up near the polarity inversion line and its end was anchored to the negative polarity field region resulting failed eruption.  Using SOT (Solar Optical Telescope) onboard {\\it Hinode}, vector magnetic field observations of NOAA AR 10953, \\inlinecite{okamoto2008} found the signature of an emerging helical flux rope beneath an active region prominence. They have found that a helical flux rope was emerging from below the photosphere into the corona along the PIL under the pre-existing prominence, and, suggested that the emergence of the helical magnetic flux rope was associated with the evolution and maintenance of the prominence as after the emergence of the flux rope prominence became more stable. In our case, we do not see any evidence of positive flux emergence near the flare and flux rope activation site. However, we have found such signature on the eastern side of PIL. Instead, a unique situation was created in the form of the activation of twisted flux rope that was going below the remnant filament channel and being failed finally. Further, the overlying filament had not changed much during the rise of the flux rope. The reconnection of this rope with the surrounding field probably led to the flare energy release at the activation site. The rise of the helical flux rope that underwent reconnection with lower coronal fields, had possibly carried material into the corona and remnant filament cavity as recently observed \\cite{okamoto2009,okamoto2010}. The rising motion of the flux rope ($\\approx$90 km s$^{-1}$) was possibly related to the outflow driven by the magnetic reconnection between the flux rope and the pre-existing coronal field in AR 11045 at the flare site.  In conclusion, we have presented here the first multi-wavelength observation of the unique event of a failed flux rope most likely due to the remnant filament in the beginning and then by an overlying coronal magnetic field. This interesting dynamics in the vicinity of rapidly emerging bipolar region and motion of its negative polarity towards the flux tube activation site may have provided the initial conditions for the flare energy build-up, and then its release due to the reconnection of the rising flux tube with the pre-existing field lines at the flare site. Further and more elaborative multi-wavelength campaign should  be carried out using the latest space ({\\it e.g.}, {\\it Hinode}, STEREO, SDO \\cite{graham2003,lemen2011}) and ground-based ({\\it e.g.}, SST, ROSA \\cite{scharmer2003,jess2010}) observatories to shed more light on such unique energy build up and release processes of solar flares associated with rapidly emerging active regions.  \\begin{acks} We thank the reviewer for his/her suggestions for our present research work. We acknowledge space missions, GOES, SOHO/MDI, {\\it Hinode}/XRT, TRACE and STEREO for providing the data used in this study.  SOHO is a project of international cooperation between ESA and NASA. {\\it 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. We are thankful to BBSO and Catania Astrophysical Observatory for providing the H$\\alpha$ data used in this study. Global High Resolution H$\\alpha$ Network is operated by the Space Weather Research Lab, New Jersey Institute of Technology. This work was supported by the Department of Science and Technology, Ministry of Science and Technology of India, by the Russian foundation for Basic Research (grants 09-02-00080 and 09-02-92626, INT/RFBR/P-38). RE acknowledges M. Ker\\'ay for patient encouragement and is also grateful to NSF, Hungary (OTKA, Ref. No. 483133) for financial support received. AKS thanks SP2RC, School of Mathematics and Statistics, The Univesity of Sheffield, U.K. for the support of collaborative visits, where the part of present research work has been carried out.  \\end{acks}"
    },
    "1107/1107.0928_arXiv.txt": {
        "abstract": "{The optically invisible infrared-source NGC~2264~IRS~1 lying north of the Cone Nebula is thought to be a massive young stellar object ($\\sim 10\\,M_{\\sun}$). Although strong infrared excess clearly shows that the central object is surrounded by large amounts of circumstellar material, no information about the spatial distribution of this circumstellar material has been available until now.} {We used the ESO Very Large Telescope Interferometer to perform long-baseline interferometric observations of NGC~2264~IRS~1  in the mid-infrared regime. Our observations resolve the circumstellar material around NGC~2264~IRS~1, provide the first direct measurement of the angular size of the mid-infrared emission, and yield direct constraints on the spatial distribution of the dust.} {We analyze  the spectrally dispersed interferometric data taken with MIDI at two different position angles and baseline lengths. We use different approaches (a geometrical model, a temperature-gradient model, and radiative transfer models) to jointly model the observed interferometric visibilities and the spectral energy distribution.} {The derived visibility values between $\\sim 0.02$ and $\\sim 0.3$ show that the mid-infrared emission is clearly resolved. The characteristic size of the MIR-emission region is $\\sim 30-60$~AU; this value is typcial for other YSOs with similar or somewhat lower luminosities. A comparison of the sizes for the two position angles shows a significant elongation of the dust distribution. Simple spherical envelope models are therefore inconistent with the data. The radiative transfer modeling of our data suggests that we observe a geometrically thin and optically thick circumstellar disk with a mass of about $0.1\\,M_{\\sun}$.} {Our modeling indicates that NGC~2264~IRS~1 is surrounded by a flat circumstellar disk that has properties similar to disks typically found around lower-mass young stellar objects. This result supports the assumption that massive young stellar objects form via accretion from circumstellar disks.} ",
        "introduction": "}The question of how massive stars ($> 8\\,M_{\\sun}$) form still remains unclear because observations of massive young stellar objects (MYSOs) are rare \\citep{Zinnecker2}. This has several reasons. First of all, these objects are surrounded by large amounts of gas and dust, which absorb nearly all the optical and near-infrared (NIR) radiation from the central source, thus making a direct observation very difficult. Second, their large distances mean that to study them we need data of high angular resolution. Other difficulties are the fast evolution of MYSOs and that they are rarely found in isolation.  Two different formation scenarios for MYSOs are currently discussed: the coalescence scenario and a formation similar to low mass protostars. In the first scenario, several low mass stars merge and finally form a single, more massive star \\citep{Bonnell1,Zinnecker1}. In the second scenario, the radiation pressure problem that would prevent spherical accretion as soon as the protostellar mass exceeds $\\sim 10\\,M_{\\sun}$ is solved by stellar outflows, which open a cavity in polar direction, through which the radiation can escape and the radiation pressure in the equatorial plane is lowered. Therefore, an infall of material in this plane is possible. Observations have provided some evidence that massive stars can form in an analogous way to low mass stars. For instance, circumstellar disks (e.g., IRAS~20126+4104, \\citealt{Cesaroni}; AFGL~490, \\citealt{Schreyer3}; IRAS~13481-6124, \\citealt{Kraus1}), molecular outflows, or collimated jets \\citep[e.g.,][]{Cunningham} are often interpreted as evidence of ongoing accretion. From the theoretical side, both the competitive accretion \\citep{Bonnell2} and the turbulent core model \\citep{McKee} agree with the formation of disks. Furthermore, simulations demonstrate that high-mass stars may form by disk accretion \\citep{Kuiper}.  Traditionally, the spatial distribution of the surrounding dust was usually studied by the modeling the spectral energy distribution (hereafter SED). Such SED model fits are highly ambiguous \\citep{Thamm,Menshchikov}. Therefore, more information is required to determine the real geometry of the circumstellar material. To test the existence of a circumstellar disk, and thus the disk accretion scenario of massive star formation, information at high angular resolution (corresponding to physical scales below 100~AU) is needed. Furthermore, only radiation at long wavelengths (e.g., infrared, sub-mm) is able to escape from the dense dusty surroundings of the star. Infrared long-baseline interferometry can thus provide direct spatial information on the required small scales.  NGC~2264 is a young, star-forming cluster with an age of about 3~Myr located in the Mon~OB~1 molecular cloud complex. Literature values for the distance of the cluster vary between 400~pc and 1000~pc, although here we adopt the most recent value of 913~pc \\citep{Baxter}. NGC~2264~IRS~1, the brightest IR-source in this region, was discovered in 1972 by D.~Allen (hence is also called Allen's Source) and has no optical counterpart. \\citet{Allen} pointed out that since NGC~2264~IRS~1 is the brightest and most luminous source in this field, it is the likely source of the radiation pressure creating the Cone Nebula. It is believed to be a young star with a mass of $9.5\\,M_{\\odot}$, spectral type B0-B2 \\citep{Thompson}, a luminosity of $3.5 \\cdot 10^3\\,L_{\\odot}$ \\citep{Harvey}, and a visual extinction of 20-30~mag. However, owing to its high extinction in the optical, all values are rather uncertain and should be interpreted with care, e.g., luminosity values between $2.3 \\cdot 10^3\\,L_{\\odot}$ \\citep{Nakano} and $4.7 \\cdot 10^3\\,L_{\\odot}$ \\citep{Schwartz} can be found in the literature.  NGC~2264~IRS~1 is associated with molecular outflows and dense molecular clumps, and in near-infrared images a twisted jet-like feature can be seen extending north of the source \\citep{Schreyer1,Ward}. The morphology of the circumstellar environment is unclear. Several studies did not find any hint of an asymmetric structure. \\citet{Schreyer2} were unable to detect any signs of a circumstellar disk in their molecular line study and \\citet{deWit} do not discern any disk-like structure in their images taken in the MIR with COMICS. However, models of a spherical symmetric distribution of the circumstellar material fail to reproduce the SED of the object in its entirety and were ruled out in several publications \\citep{Tokunaga,Harvey}. However, the spatial resolution of all these studies may not be sufficient to resolve a disk-like structure.  In this paper, we present mid-infrared interferometric observations of the MYSO NGC~2264~IRS~1 performed with the Mid-Infrared Interferometric Instrument (MIDI) at the Very Large Telescope Interferometer (VLTI). The MIDI data and their reduction are shown in Sect.~\\ref{Observations}. Different modeling approaches (geometrical, temperature-gradient, and radiative transfer models) for the visibilities and the SED are presented in Sect.~\\ref{Modelling}. Implications for the object and final conclusions are discussed in Sect.~\\ref{Discussion}. ",
        "conclusions": "} We have presented MIDI mid-infrared interferometric observations of the massive young stellar object NGC~2264~IRS~1. The observed visibilities provide \\textit{no} hint of multiplicity of the central source in the separation range from $\\sim 30$~AU (resolution limit) to $\\sim 230$~AU (slit width limit). Geometric models (Sect.~\\ref{Geometry}) suggest a size of $\\sim 30-60$~AU for the mid-infrared emission region. This value is in good agreement with the mid-infrared sizes of similar objects.  Our modeling shows that a simple temperature-gradient model for a circumstellar disk can reproduce the SED at mid- and far-IR wavelengths, but fails to reproduce the SED in the near-IR and the observed visibilities; this model can therefore be ruled out.  We have performed a detailed radiative transfer modeling of the observed SED and visibilities with the RADMC code. It suggests a scenario of a geometrically flat but optically thick circumstellar dust disk. We also used the fitting tool by \\citet{Robitaille}, which leads to similar disk models and therefore confirms the results obtained with RADMC, whereas envelope models found within the Robitaille grid did not provide similarly good fits to SED and visibilities simultaneously. Although there is no observational evidence, the possible presence of a large ($r \\geq 400$~AU) spherical envelope around the disk cannot be excluded by the data. The mass of the circumstellar disk is about $0.1\\,M_\\odot$. The data suggest an inclination angle of $\\sim 30\\degr$ and a position angle of $\\sim 40\\degr$ for the orientation of the disk. These values are consistent with an overall geometrical model based on the jet-like feature seen to the north-east of IRS~1. Comparing NGC~2264~IRS~1 to other YSOs, the size of its MIR emitting region seems to be typical of its luminosity and mass. However, the large scatter of sizes in this range of luminosities and masses points towards a wide variety of (disk) morphologies among these objects.  More observational research, such as future MIR interferometric observations, will provide tighter constraints on the circumstellar material around this interesting source and therefore will help us to  clarify our understanding of the disks around high-mass YSOs.  \\begin{acknowledgement} We would like to thank the anonymous referee for his suggestions and comments that helped to improve the paper. We would also like to thank K.~Dullemond for explanations and discussions about RADMC. We gratefully acknowledge funding of this work by the German \\textit{Deutsche Forschungsgemeinschaft, DFG} project number PR~569/8-1. \\end{acknowledgement}"
    },
    "1107/1107.5020_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  The deviation from perfect Gaussian statistics of the primordial fluctuations promises to allow a more precise discrimination between competing scenarios of the early Universe than has hitherto been possible. A central object in this respect is the primordial bispectrum: the three-point correlation function of the primordial curvature perturbation $\\z$ in Fourier space $\\langle \\z (\\boldsymbol{k}_1)  \\z (\\boldsymbol{k}_2)  \\z (\\boldsymbol{k}_3)\\rangle$. An easy way to acknowledge the wealth of information contained in this object is to note that it is a function of three variables, contrary to the power spectrum which is a function of only one variable. Consequently, considerable amount of attention has been given to understand and classify the bispectrum's dependence on the configuration of the momenta $\\boldsymbol{k}_1,\\boldsymbol{k}_2,\\boldsymbol{k}_3$, what is called its shape \\cite{Babich:2004gb,Fergusson:2008ra}. Beyond the standard local \\cite{Komatsu:2001rj} and equilateral \\cite{Creminelli:2005hu} shapes of non-Gaussianities, lots of other patterns, corresponding to different mechanisms for generating the primordial fluctuations, have now been identified (see for instance \\cite{Chen:2010xk} and \\cite{Liguori:2010hx} for theoretical and observational reviews respectively). In practice, shapes are derived from operators in the third-order action for the fluctuations, so that a natural question to ask is: what is the set of such possible cubic operators?  In this respect, we elaborate in this short note on a technical point whose important consequences have not yet been fully acknowledged: some cubic operators that are \\textit{a priori} independent, \\textit{i.e.} not simply related by integrations by part, can actually be related to each other by using the linear equations of motion, a procedure which has recently been stressed to be valid for computing correlation functions \\cite{Arroja:2011yj,Burrage:2011hd}. This considerably limits the range of truly independent cubic operators, and hence of possible shapes of non-Gaussianities.  We demonstrate this in the framework of Hordenski's most general single-field scalar-tensor theory with second-order equations of motion \\cite{Hordenski}, which has been ``resurrected'' recently  \\cite{Charmousis:2011bf} and which has been shown to be equivalent to generalized Galileons in four dimensions \\cite{Deffayet:2011gz} in reference \\cite{Kobayashi:2011nu}. The corresponding third-order action has been calculated in references \\cite{Gao:2011qe,DeFelice:2011uc}, where the appearance of three cubic operators absent in the simpler class of models known as $k$-inflation was stressed. As we pointed out to Gao and Steer for a revised version of their preprint \\cite{Gao:2011qe}, and which we now explain in further details, these ``new'' operators can actually be expressed in terms of ones which have already been studied in depth in the context of $k$-inflation \\cite{Burrage:2011hd}, thereby alleviating the need for their separate study.  In the next section, we quote the second-and third-order action for scalar perturbations in the most general single-field scalar-tensor theory with second-order equations of motion and we explain the strategy behind our explicit calculations in section \\ref{Calculations}. We briefly conclude in section \\ref{Conclusion}. ",
        "conclusions": "} \\label{Conclusion}  In this short note, we have reconsidered the status of the operators that appeared in references \\cite{Gao:2011qe,DeFelice:2011uc} in the third-order scalar action of the most general scalar-tensor theory with second-order equations of motion. There, three operators were identified besides the five ones that naturally appear in the simpler class of models known as $k$-inflation and whose properties are well known. Emphasizing that it is valid for computing correlation functions to use the linear equation of motion to simplify the third-order action, we have shown that these three operators are actually redundant and can be exactly traded for the five well known ones.  Under the slow-varying approximation and in the Bunch-Davies vacuum, these five independent cubic operators generate primordial bispectrum shapes either of equilateral type -- for $\\dot \\z^3$, $\\dot \\zeta \\partial_i \\zeta \\partial^i \\left( \\p^{-2} \\dot \\z \\right)$ and $\\partial^{2}\\zeta  \\left(\\partial_i \\p^{-2} \\dot \\z  \\right)^{2}$ -- or of local type -- for $\\z \\dot \\z^2$ and $\\z (\\p \\z)^2$. However, even in these circumstances, we should point out that this does not entail that nothing else can be expected from single field models of inflation. For instance, the orthogonal shape of non-Gaussianities \\cite{Senatore:2009gt}, which has never been shown to be generated in a particular $k$-inflationary model, was recently shown to emerge in a simple model of DBI Galileon inflation \\cite{RenauxPetel:2011dv,RenauxPetel:2011uk}. The set of possible shapes of course gets even larger when features or/and non-Bunch Davies vacuum are considered (see for example \\cite{Chen:2010bk} for a recent illustration).   Eventually, let us note that Eq. \\refeq{3} expresses the fact that the operator $\\dot \\z (\\p \\z)^2$ -- which typically generates a primordial bispectrum predominantly correlated with the equilateral template -- can be replaced by a combination of the three operators $\\dot \\z^3$, $\\z \\dot \\z^2$ and $\\z (\\p \\z)^2$, the last two generating a bispectrum whose shape is predominantly correlated with the local template. This -- and more generally the fact that cubic operators that are \\textit{a priori} independent can actually be related on using the linear equations of motion -- indicates that one should be cautious with inferring results for the global shape of the primordial bispectrum from the type of operators present in a given form of the third-order action. \\\\  \\noindent  {\\bf  Acknowledgments:} I would like to thank David Seery for useful comments on a draft version of this paper. I am supported by the STFC grant ST/F002998/1 and the Centre for Theoretical Cosmology."
    },
    "1107/1107.2436_arXiv.txt": {
        "abstract": "We show that simple 2 and 3-layer flux-transport dynamos, when forced at the top by a poloidal source term, can produce a widely varying amplitude of toroidal field at the bottom, depending on how close the meridional flow speed of the  bottom layer is to the propagation speed of the forcing applied above the top layer, and how close the amplitude of the $\\alpha$-effect is to two values that give rise to a resonant response. This effect should be present in this class of dynamo model no matter how many layers are included. This result could have implications for the prediction of future solar cycles from the surface magnetic fields of prior cycles. It could be looked for in flux-transport dynamos that are more realistic for the Sun, done in spherical geometry with differential rotation, meridional flow and $\\alpha$-effect that vary with latitude and time as well as radius. Because of these variations, if resonance occurs, it should be more localized in time, latitude and radius. ",
        "introduction": "Dikpati et al (2006) first used a flux transport dynamo calibrated to the Sun (Dikpati et al 2004) to simulate and predict solar cycle peaks from the record of past surface magnetic field patterns. This was done mathematically by forcing the dynamo equations at the top boundary, with a forcing function derived from past surface magnetic fields. Flux transport dynamos, and indeed all dynamos, have their own unforced, usually complex frequencies of excitation that are commonly found by treating the dynamo equations as an eigenvalue problem. Many naturally occurring and man-made systems have such properties.  When a physical system that has natural freqencies is excited by external forcing whose own frequency is close to one of the natural ones, there can be resonance produced--that is, the system will be excited strongly by the forcing compared to the case where the forcing frequency is not close to a natural one. The purpose of this paper is to explore the possibility of resonance in flux-transport dynamos relevant to the solar cycle.  In flux transport dynamos, there are several physical properties that help determine the unforced frequencies of the system. These include differential rotation, meridional circulation, the so-called $\\alpha$-effect, or kinetic helicity, and turbulent magnetic diffusion. It is now well established (Dikpati and Charbonneau, 1999) that unless the magnetic diffusivity is very large, meridional flow at the bottom of the dynamo layer is primarily responsible for the real part of the natural frequency of the dynamo, which determines the speed with which induced toroidal and poloidal fields near the bottom migrate toward the equator. Therefore the closeness of the frequency of forcing at the top to the speed of the flow at the bottom could help determine how much dynamo response there is.  Since the forcing at the top is created by emergence of concentrated magnetic flux from the bottom, in the form of active regions, and the rate of movement of the zone where active regions are found moves toward the equator (not coincidentally) at a rate close to the meridional flow speed near the bottom, we might expect the conditions for resonance to occur in the bottom layer to be favorable. On the other hand, we know from observations (Ulrich, 2010 and references therein) that the meridional flow at the top of the convection zone is toward the poles, opposite to the propagation of the surface forcing as well as 5-10 times faster. Thus we should not expect resonance to occur near the surface.  It is also well known (Ulrich 2010 and references therein) that the meridional circulation varies with time. This time variation is now being incorporated into a flux-transport dynamo used for prediction by Dikpati and colleagues. In the 2006 prediction, meridional circulation generally was kept fixed in time. Dikpati et al (2006), Dikpati and Gilman (2006) recognized that such time variations could be important, but felt they lacked sufficient knowledge of its variations to include them. They adjusted the time-independent meridional flow amplitude to give the average period of the past solar cycles, and stretched or compressed all the surface forcing data to the same period, to avoid any artificial or non-physical mismatches between the natural dynamo period and the period of the forcing.  But there can also in principle in the Sun be real differences between the period of the top forcing that was created by the previous cycle, and the freqency of equatorward propagation associated with the meridional flow speed at the bottom. In dynamos forced at the top with a specified period, the amplitude of the induced fields within the dynamo domain will be affected by this frequency difference. The model we present here in effect studies how this amplitude is affected, by treating the meridional flow at the bottom as a free parameter while keeping the frequency of the top forcing fixed.  In the real sun, the cycle period varies from cycle to cycle, as does the speed of the meridional flow and its profile with latitude. Ultimately it is highly desirable to include both such variations. This can be done by use of data assimilation techniques applied to both the surface forcing and meridional flow variations. As we said above, Dikpati and colleagues are doing that now. When that is accomplished, they may find that resonance plays some role. In this paper, we anticipate that possibility and focus on possible resonances by using a much simpler dynamo model than used in Dikpati and Gilman (2006), namely one that has no more than two three layers in the radial direction.  Such an approach has the advantage of speed while retaining important physical processes. But such a simple model would have little value as a tool for prediction, because it could not be calibrated well in detail to the sun, since it would have few degrees of freedom. It also may overestimate the importance of resonance for the same reason. The cautions expressed in Roald (1998) about the limits of dynamo models with one or two layers are well taken. Nevertheless, since the forced dynamo problem has only begun to be studied, particularly in the solar case, using a really simple model initially may give useful guidance about what to look for with a more realistic version. It is in this spirit that we report on these calculations here.  Resonance has been studied in dynamos previously, but the literature is small. General examples include Strauss (1986) and Reshetnyak (2010). Resonance in the geodynamo has been studied by Stefani and Gerberth (2005) and Fischer et al (2008). Studies for disks and galaxies include Chiba (1991), Schmitt and R\\\"udiger (1992), Kuzanyan and Sokoloff (1993), and Moss (1996). We have not located any previous studies specific to the Sun in which resonance has been explicitly identified and highlighted. However in all these areas there are almost certainly model studies in which some form of resonance is playing a role, but has not been brought out in the analysis of results. Any dynamo in which inputs such as flow fields or turbulent parameters such as the $\\alpha$-effect are allowed to vary with time, either imposed or by nonlinearities internal to the system, such as 'quenching', could display behavior related to resonance. ",
        "conclusions": "Our principal conclusion is that, at least in this simple dynamo model, resonance will always be found, provided the right parameter values are chosen. Furthermore, it occurs for parameter choices that are plausible for the sun. It occurs in the lowest layer of the model, where the meridional flow toward the equator is most likely to match, or nearly match, the propagation speed of the forcing imposed at the top, corresponding to the photosphere on the sun. And the effect is large--a factor of 10-100 amplification of toroidal field compared to the case where the parameter values are far from the ones for which resonance is predicted. We acknowledge that this large a difference in amplitude for different parameter values is partly due to the model being kinematic. If $jXB$ forces were included, the peaks would almost surely be smaller.  The solar convection zone is a spherical shell, not an infinite layer, so how well should we expect our results to apply to the Sun? In our model, the differential rotation $s$ is a constant, so the rotation is a linear function of 'latitude' only, and independent of depth. In the Sun, differential rotation varies with both latitude and depth. In our model, meridional flow is independent of latitude, and there is no vertical motion, while in the Sun, meridional flow is a closed circulation confined by the boundaries of the convection zone as well as the poles. Meridional circulation is also known to vary with time on a variety of timescales. The $\\alpha$-effect will also be a function of latitude as well as time.  These differences between our model and the Sun imply that resonance, if it occurs, would likely be found in more localized regions of the dynamo domain. But in principle it could still occur. If a particular location experiences resonance or near resonance for some period of time, then the toroidal field might amplify quickly there, leading to locally anomalously large field. How quickly presumably depends on the nearness to conditions for resonance, and how long these conditions last. Could such a sudden amplification be a precursor for the formation of a rising flux tube that creates a new active region? This seems like a hypothesis worth testing in the future.  It is also possible that in the convection zone conditions for resonance could persist for large fractions of a solar cycle and perhaps even from one cycle to the next. If this happens, it could contribute to the evolution of the envelope of cycle amplitudes. On extremely long time scales could persistent conditions favoring resonance be responsible for 'grand maxima', and conditions unfavorable for resonance lead to 'grand minima', such as the Maunder minimum? Clearly these are very speculative questions, requiring much more sophisticated solar dynamos than we have used here, to answer.  In flux-transport dynamos in spherical shells that are forced from the top, it is the radial flow of the closed meridional circulation that transports the forcing signal including its frequency to the bottom. In our infinite plane cartesian model, there is no radial flow, so the signal should be thought of as getting to the bottom locally via vertical diffusion and induction. We can think of the meridional circulation in the upper layer closing at infinity and returning in the lower layer.  Since we know that not all sunspot cycles have the same duration, the meridional flow will be bringing a frequency to the bottom that could be different than that implied by the meridional circulation there at the time the active regions were formed, that gave rise to the forcing frequency. Also, a change in the speed of the flow carrying the forcing signal will change the frequency of that signal by Doppler-shifting it. But in percentage terms the differences will not be large, so there could be continually evolving 'nearness' to conditions needed for resonance, creating continuous changes in cycle amplitude, in addition to the usual evolution of a sunspot cyle through ascending, peak, declining and minimum phases. Again, investigating such possibilities requires much more realistic dynamo models than we have used here.  THe National Center for Atmospheric Research is sponsored by the National Science Foundation. This work is partially supported by the NASA Living with a Star Program through Award NNX08AQ34G."
    },
    "1107/1107.0269_arXiv.txt": {
        "abstract": "{% In this study we determined precise orbital and physical parameters of the very short period low-mass contact binary system CC Com. The parameters are obtained by analysis of the new CCD data with the archival spectroscopic data. The physical parameters of the components derived as $M_\\textrm{c}$ = 0.717(14) $M_{\\odot}$, $M_\\textrm{h}$ = 0.378(8) $M_{\\odot}$, $R_\\textrm{c}$ = 0.708(12) $R_{\\odot}$, $R_\\textrm{h}$ = 0.530(10) $R_{\\odot}$, $L_\\textrm{c}$ = 0.138(12) $L_{\\odot}$, $L_\\textrm{h}$ = 0.085(7) $L_{\\odot}$, and the distance of the system is estimated as 64(4) pc. The times of minima obtained in this study and with those published before enable us to calculate the mass transfer rate between the components which is  $1.6\\times10^{-8}$ M$_{\\odot}$yr$^{-1}$. Finally, we discuss the possible evolutionary scenario of CC Com.} ",
        "introduction": "One of the crucial parameter that determines the evolutionary stages of a binary is its orbital parameter. Because of their unusual behaviours short period systems like CC Com, GSC 1387-0475 (Yang et al. 2009) and V523 Cas (K\\\"{o}se et al. 2009) are important in evolutionary studies.  CC Com was discovered by Hoffmeister (1964) and has been intensively studied photometrically by Rucinski (1976),  Breinhorst \\& Hoffmann (1982), Bradstreet (1985), Zhou (1988), Linnell \\& Olson (1989) and Zola et al. (2010). Over the last four decades an intense spectroscopic studies were made by Rucinski et al. (1977), McLean \\& Hilditch (1983), and Pribulla et al. (2007). In this studies spectroscopic mass ratio was found as 0.52(3), 0.47(4), and 0.53(1), respectively. Photometric mass ratio, on the other hand, was given as 0.59 by Zhou (1988) and 0.51 by Linnell \\& Olson (1989). The physical parameters of the components have not been measured by simultaneous analysis of spectroscopic and photometric data.   Orbital period study of weakly contact binary CC Com has been the subject of many papers. These studies indicate a systematic period decrease. In the literature the period variation ($dP/dt$) of CC Com was given as  $-4.4\\times10^{-8}$ d~yr$^{-1}$, $-4.0\\times10^{-8}$ d~yr$^{-1}$, and $-2.0\\times10^{-8}$ d~yr$^{-1}$ by Qian (2001a), Yang \\& Liu (2003), and Yang et al. (2009), respectively. Yang et al. (2009) presented cyclic variations superimposed on a parabolic period variation and discussed this feature as an indication of a third body or stellar activity.  Contact model solution has been suggested by Rucinski (1976) and related UBV parameters have been given in that study. Breinhorst \\& Hoffmann (1982) studied the effects that can cause variations in minima depths assuming no period change. The filling factor variation has been discussed by Linnell \\& Olson (1989) based on their $u$, $y$ and $I$ photometric data analysis. Recently, Zola et al. (2010) gave the photometric elements of the system by using Wilson-Devinney code.  In this paper, we present photometric analysis and orbital period study of the short-period eclipsing binary system CC Com. First, the new observations with new linear ephemeris are presented. Then the mass transfer rate between the components is estimated by period study. Next, the light curve of the system is modeled by comparing the results with the results of previous works. Finally, obtained results are presented with a discussion of possible evolutionary stages of the binary. ",
        "conclusions": "\\begin{table} \\scriptsize \\begin{center} \\caption{Absolute parameters of CC Com. The standard errors 1$\\sigma$ in the last digit are given in parentheses. HC denotes the hot component and CC stand for the cool component.}\\label{Tab:physpar:CCCom} \\begin{tabular}{llll} \\hline Parameter                                        &Unit                      & CC          & HC   \\\\ \\hline Mass (\\textbf{M})                                         &$\\rm{M_{\\odot}}$      & $0.717(14)$         & $0.377(8)$      \\\\ Radius (\\textbf{R})                                       &$\\rm{R_{\\odot}}$      & $0.708(12)$         & $0.530(10)$      \\\\ Temperature (\\textbf{T$_{\\rm eff}$})                      &$\\rm{K}$              & $4300$              & $4200(180)$    \\\\ Luminosity (\\textbf{L})                                   &$\\rm{L_{\\odot}} $     & $0.138(12)$         & $0.085(7)$      \\\\ Surface gravity $(\\log\\,\\textbf{g})$                      &cgs                   & $4.59$              & $4.57$          \\\\ Absolute bol. mag. (\\textbf{M$_b$})            &mag                   & 6.90(14)            & 7.43(18)            \\\\ Absolute vis. mag. (\\textbf{M$_V$})                &mag                   & 7.86(16)            & 8.25(20)           \\\\ Period change rate ($\\dot{\\textbf{P}}$)                   &d/yr                  &\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,$-1.34\\times10^{-8}$ &      \\\\ Mass transfer ratio ($\\dot{\\textbf{M}}$)                  &M$_\\odot$/yr          &\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,$-1.6\\times10^{-8}$ &      \\\\ Distance (\\textbf{d})                                     &pc                    &\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,\\,64(4)         &      \\\\ \\hline \\end{tabular} \\end{center} \\end{table} Period variation and light-curve analysis of the low temperature contact binary (LTCB) system CC Com have been studied. The physical parameters of the system have been determined with the simultaneous solution of our two colour light curves (\\emph{VR}) and the spectroscopic study of Pribulla et al. (2007).   The measured physical parameters of CC Com have been presented in Table~\\ref{Tab:physpar:CCCom} with their errors. It seems that there is a good agreement between the derived physical parameters of CC Com and other LTCBs (Yakut \\& Eggleton, 2005). We estimate the distance to CC Com by using the results obtained from radial velocity and light curve analysis. For that purpose the total brightness (V=11\\fm30) and light ratio of the components have been used. The results indicate a distance of 65 pc and 63 pc for the hot and cool components, respectively. The mean of these estimates gives the distance of the system as 64(4) pc. This value is 20\\% smaller than the value given by SIMBAD database.   We have collected and analyzed the times of minima for CC Com. The (O--C) diagram shows a downward parabola. This property can be explained as a mass transfer from the more massive component to the less massive one. The quadratic term of Eq.~\\ref{cccomEq2} shows that the orbital period of the system decreases at a rate of ${dP}/{dt}=1.34\\times 10^{-8}$\\,days/yr as a result of mass transfer rate of $1.6\\times 10^{-8}$ M$_\\odot$/yr. In this study, the residuals of the period variation do not show any reliable sine-like variation (Fig.~\\ref{Fig:CCCom:OC}b). Assuming the existence of the third body we solved the data, however, because of the poor data it is hard to find any evidence for a third body in the system. On the other hand, Yang et al. (2009) assumed a third body in the system; solved the residuals and determined the third orbit parameters.   Most of the contact binaries show the O'Connell effect in their light curves. This effect is due to large cool starspots (see for details Kalomeni et al. 2007). The variation seen in Fig.~\\ref{Fig:CCCom:OC}b can also be explained by stellar activity. Similar residuals have been observed in many contact binary systems (e.g., XY Leo, Yakut et al. 2003). The light curve solution indicates that 6\\% of the primary star's surface is covered by a cold spot. In this case, Applegate mechanism (Applegate 1992) can be responsible from the variation in the orbital period and the occurrence of the non-periodic change.   Different scenarios have been proposed for the evolution and structure of contact binaries. The prevailing theory among them is the thermal relaxation oscillation (TRO) theory, proposed by Lucy (1976), Flannery (1976),  Robertson \\& Eggleton (1977), Yakut \\& Eggleton (2005). The TRO model can explain successfully  the evolution, structure and the observed properties of contact binary systems (Wang, 1994, Qian 2001b, Van Hamme, 2001, Webbing 2003, Paczy{\\'n}ski et al. 2006, Zhu et al. 2010).  Stepien (2006), also proposed a scenario explaining the evolution of contact binary systems. In short period binaries like CC Com, angular momentum loss plays a crucial role in their evolution. For a close detached binary system with initial parameters 1.19 $\\rm{M_{\\odot}}$ + 0.94 $\\rm{M_{\\odot}}$ and  0.75 days, the mass of the system is estimated to decrease by $\\sim 15\\%$  (0.97 $\\rm{M_{\\odot}}$+0.83 $\\rm{M_{\\odot}}$) until \\textsc{RLOF} at P = 0.31 days, and it reaches the contact phase at 0.88 $\\rm{M_{\\odot}}$ + 0.91 $\\rm{M_{\\odot}}$, 0.28 days. This scenario can explain the evolutionary stages of CC Com with TRO mode like the other LTCBs that are discussed in detail in Yakut \\& Eggleton (2005).   Additionally, the evolution of short period binaries is also important in measuring the gravitational waves. Binary systems with very short period can form gravitational waves in detectable range (Ju et al., 2000, K{\\\" o}se \\& Yakut, 2011). We have estimated the amplitude of the gravitational wave in CC Com binary system as ($\\log (\\textrm{h})$) -20.6. Hence, the system CC Com is an important source for interferometers and its amplitude is within the detection limit of detectors such as LISA that can detect accurate gravitational waves. The relatively close distance of the system makes it also important target to study."
    },
    "1107/1107.4483_arXiv.txt": {
        "abstract": "In the internal shock model of gamma ray bursts ultrahigh energy muons, pions, neutrons and kaons are likely to be produced in the interactions of shock accelerated relativistic protons with low energy photons (KeV-MeV). These particles subsequently decay to high energy neutrinos/antineutrinos and other secondaries. In the high internal magnetic fields of gamma ray bursts, the ultrahigh energy charged particles ($\\mu^+$, $\\pi^+$, $K^+$) lose energy significantly due to synchrotron radiations before decaying into secondary high energy neutrinos and antineutrinos. The relativistic neutrons decay to high energy antineutrinos, protons and electrons. We have calculated the total neutrino flux (neutrino and antineutrino) considering the decay channels of ultrahigh energy muons, pions, neutrons and kaons. We have shown that the total neutrino flux generated in neutron decay can be higher than that produced in $\\mu^+$ and $\\pi^+$ decay. The charged kaons being heavier than pions, lose energy slowly and their secondary total neutrino flux is more than that from muons and pions at very high energy. Our detailed calculations on secondary particle production in $p\\gamma$ interactions give the total neutrino fluxes and their flavour ratios expected on earth. Depending on the values of the parameters (luminosity, Lorentz factor, variability time, spectral indices and break energy in the photon spectrum) of a gamma ray burst the contributions to the total neutrino flux from the decay of different particles (muon, pion, neutron and kaon) may vary and they would also be reflected on the neutrino flavour ratios. ",
        "introduction": "Extensive airshower arrays \\cite{pa1,pa2,hires,agasa,yakutsk} have detected a large number of ultrahigh energy cosmic ray events. The origin of these cosmic rays are yet to be identified. They may come from supernova remnants (SNRs), active galactic nuclei (AGN), gamma ray bursts (GRBs) or some unknown sources. Some of the ultrahigh energy cosmic rays (protons and nuclei) may interact with low energy radiations and matter inside the source producing secondary high energy photons and neutrinos. Protons can be shock accelerated to $\\sim 10^{21}$ eV inside GRBs by Fermi mechanism. Shock accelerated protons may loose energy by synchrotron cooling and $p\\gamma$ interaction inside the source depending on the magnetic field and low energy photon density respectively. The high energy neutrino flux produced in $p\\gamma$ interactions have been calculated in detail in many earlier papers \\cite{vietri,wax1,rach,alv,nay1,guetta,bhat,nay2,murase1,becker1,wax2,murase2}. It has been noted earlier that $pp$ interactions are only important for photospheric radii of GRB fireballs \\cite{murase3,dai}. The interactions of shock accelerated protons with low energy radiations ($p\\gamma$) lead to the production of ultrahigh energy mesons and leptons $p\\gamma\\rightarrow{\\pi^{+,0}}X,\\, {\\pi^{+}} \\, \\rightarrow \\, {\\mu^{+}}\\nu_{\\mu}, \\, {\\mu^{+}} \\, \\rightarrow \\, {e^{+}} \\bar\\nu_{\\mu}\\nu_e$, $X$ can be neutrons and they will decay to protons, electrons and antineutrinos of electron flavour ($n\\,\\rightarrow p +e^-+\\overline{\\nu}_e$) in 886 sec in their rest frames. The other secondary products in $p\\gamma$ interactions are $p\\gamma \\rightarrow{K^{+,0}}X$ where X can be either $\\Lambda^0$, $\\Sigma^0$ or $\\Sigma^+$. Kaons decay to lighter mesons, leptons and neutrinos. If the magnetic field inside the GRBs is very high then the ultrahigh energy muons, pions and kaons lose energy significantly before decaying to neutrinos. Synchrotron cooling of charged muons and mesons would lead to decrease in the fluxes of neutrinos from their decay. Although, the cross-section for kaon production from p$\\gamma$ is much less than that for photo-pion production, kaons being heavier cool at a much slower rate compared to muons and pions. As a result at very high energy the kaon decay channel of neutrino production becomes significant compared to the pion decay channel \\cite{and,asan,wal1,wal2,kal1}. Shock accelerated protons are expected to produce high energy neutrons in various interactions, and these unstable particles subsequently decay to leptons and protons.  In the context of UHECR production from GRBs the neutron decay channel was considered earlier in \\cite{der}. The energy and time dependence of neutron, neutrino fluxes from GRBs were calculated within the scenario of external shock model in this paper. Also the radiation halos from neutron decay electrons, protons and their detectability in different wavelengths were studied.  In a recent paper \\cite{ahl} the authors have explored the possiblity of detecting high energy diffuse neutrino flux from GRBs by IceCube, assuming the diffuse UHECR flux as detected by HiRes is the proton flux produced by the decay of cosmic ray neutrons in GRB fireballs.  In our work we have not assumed that GRBs are the sources of the UHECR events observed by Pierre Auger \\cite{pa1,pa2}, HiRes \\cite{hires} or AGASA \\cite{agasa} experiments. We have considered the possibility of cosmic ray acceleration in individual GRBs to ultra high energies to explore the parameter dependence of the high energy neutrino fluxes produced in decay of various particles. We discussed the importance of the neutron decay channel in our earlier paper \\cite{ree} for cosmic accelerators with high internal magnetic fields. In this paper we have calculated the neutrino fluxes from GRBs in the internal shock model through muon, photo-pion, neutron and kaon decay, including the effect of proton energy loss by synchrotron cooling and $p\\gamma$ interactions, synchrotron cooling of charged muons, pions, and kaons in the internal magnetic field of GRB. After including proton energy losses, we find that the neutron decay channel of high energy antineutrino production may remain important in GRBs depending on the values of the GRB parameters. Moreover, the various decay processes may show distinct features in the neutrino flavour ratios.  Due to the ongoing experimental activities to detect high energy neutrinos from GRBs \\cite{abbasi0,abbasi1,abbasi2,abbasi3,abbasi4} this field has remained exciting. IceCube collaboration has claimed to reach the sensitivity of detecting neutrino flux from GRBs at TeV energy. IceCube opearted in a 40-string configuration from April 5, 2008 to May 20, 2009 \\cite{abbasi1}. They considered 117 bursts and no events were detected above the atmospheric background. For cosmic neutrinos with an $E^{-2}$ energy spectrum an integral flux limit of $E^2\\phi\\leq 3.6\\times 10^{-8} GeV cm^{-2} sec^{-1}sr^{-1}$ has been found in the energy interval of $2\\times10^6-6.3\\times10^{9} GeV$ \\cite{abbasi2}. In particular the non-detection of neutrinos from GRBs \\cite{abbasi4} has placed a tighter upper bound which is 3.7 times below the theoretical predictions \\cite{wax1,rach, guetta, ahl} combining 40 and 59 strings. IceCube collaboration has concluded \\cite{abbasi4} that GRBs are not the only sources of cosmic rays above energy 1 EeV or the efficiency of neutrino production is much lower than the current predictions. The values of the bulk Lorentz factor $\\Gamma$ and the ratio of energies in protons to electrons have been constrained at the $90\\%$ confidence level in their Fig.4. H\\\"ummer et al. \\cite{hum} have recalculated the neutrino fluxes from GRBs and concluded that their result is significantly below the limit set by IceCube experiment operated in 40 strings. It has been pointed out by Li \\cite{Li} that the theoretical prediction of neutrino flux from GRBs in IceCube papers is an overestimation, and the non-detection of GRB neutrinos is consistent with the correct theoretical estimation of neutrino flux. He et al. \\cite{He} have calculated the neutrino flux from GRBs including proton energy losses and they have also suggested that the neutrino flux predicted theoretically in the papers by IceCube collaborations is an overestimation. In calculating the photon number density IceCube collaboration approximated the energy of all photons by the break energy of the photon spectrum. The recalculated neutrino flux by He et al. \\cite{He} from 215 GRBs observed by IceCube 40 and 59 string configurations is about $36\\%$ of the $90\\%$ confidence upper limit obtained by IceCube collaboration and it is consistent with non-detection of GRB neutrinos by IceCube detector. The objective of the present work is to show that the relative importance of the different channels (pion, muon, neutron and kaon decay) of neutrino production in GRBs depends on the values of the GRB parameters. We have compared our calculated flavour ratios with earlier work \\cite{wax3,wal2}. After including muon energy loss and kaon decay channels in detail at high energy, our flavour ratios are lower than the constant value of $\\phi_{\\nu_{\\mu}}/(\\phi_{\\nu_{e}}+ \\phi_{\\nu_{\\tau}})=0.64$ found in earlier work. ",
        "conclusions": "Gamma ray bursts have been speculated to produce high energy particles like cosmic rays, gamma rays and neutrinos. Ultrahigh energy muons, pions, neutrons and kaons could be produced in GRBs due to p$\\gamma$ interactions. These particles subsequently decay to produce very high energy neutrinos and antineutrinos. Charged particles lose their energy due to synchrotron cooling in the magnetic field before decaying to lighter particles. Kaons suffer synchrotron loss at higher energy compared to muons and pions. As a result at higher energy the kaon decay channel of neutrino production dominates over the muon and pion decay channels.We have included all the decay channels of charged ($K^+$) and neutral kaons ($K^0_L$ and $K^0_S$) to study this effect. The total neutrino flux from neutron decay channel may be more important than that from muon and pion decay channels depending on the values of the parameters of GRBs. The flavour ratios of neutrinos show distinct features depending on the values of the GRB parameters and the dominance of different decay processes at different energies. IceCube has set limit on the GRB neutrino flux. Future detectors with higher sensitivities will be able to reveal the role of GRBs as a UHECR accelerator."
    },
    "1107/1107.5759_arXiv.txt": {
        "abstract": "The small community of the Tom\\'ar\\^aho, an ethnic group culturally derived from the Zamucos, have become known in the South American and world anthropological scenario in recent times. This group, far from the banks of the Paraguay river, remained concealed from organized modern societies for many years. Like any other groups of people in close contact with nature, the Tom\\'ar\\^aho developed a profound and rich world-view which parallels other more widely researched aboriginal cultures as well as showing distinctive features of their own. This is also apparent in their imagery of the sky and of the characters that are closely connected with the celestial sphere.  This paper is based on the lengthy anthropological studies of G. Sequera. We have recently undertaken a project to carry out a detailed analysis of the different astronomical elements present in the imagined sky of the Tom\\'ar\\^aho and other Chamacoco ethnic groups.  We will briefly review some aspects of this aboriginal culture: places where they live, regions of influence in the past, their linguistic family, their living habits and how the advancement of civilization affected their culture and survival. We will later mention the fieldwork carried out for decades and some of the existing studies and publications. We will also make a brief description of the methodology of this work and special anthropological practices. Last but not least, we will focus on the Tom\\'ar\\^aho conception of the sky and describe the research work we have been doing in recent times. ",
        "introduction": "The act of living in a certain culture may be displayed in many different ways. One of these is creativity, at stake when observing the changing phenomena of the surrounding natural world: giving names to things; naming people, animals and plants; in short, building up the knowledge and symbols to enable them to acquire an identity and survive as a society. The present-day groups aspire to this, completely paralleling the cultures in the past, for whom this was a strong motivation---both of the more technologically developed groups as well as the ones who did not follow the frantic development of `civilization' and remained in closer contact with nature. The small community of the Tom\\'ar\\^aho is an example of the latter.  A first contact with members of this ethnic group took place in 1986 when one of the authors (Sequera) was able to see for himself the situation of near slavery and strong dependence the members of this people had to put up with in their relationship with the Carlos Casado company, which was industrializing the region. This Anglo-Argentinian enterprise, in complicity with the Paraguayan government since the end of the 19th century (a few years after the end of the War of the Triple Alliance which weakened the nation considerably), had illegally taken the lands of the Tom\\'ar\\^aho and exploited the quebracho woods extensively in order to obtain timber for railway sleepers and to stock their tanning plants. With no recognized rights of any kind, the natives were made to work as axemen under hard labour conditions and to be, in their turn, witnesses to the plunder of the Chaco woods. In this way they became helpless victims of the subjugation of the environment that was vital to them (Sequera 2006).  Given the importance of carrying out an anthropological study of this ethnic minority, comparative studies were undertaken of ethnographic sources in different libraries and research centres. These have helped to clarify our view of the possible origin of the Tom\\'ar\\^aho and uncovered mentions of them in Jesuits' letters ({\\it cartas annuas}) and old chronicles. Sequera, moreover, lived among the natives for several periods between 1987 and 1992, which enabled him to carry out a methodical programme of recognition and transcription of the Tom\\'ar\\^aho language, to compile a detailed inventory of the social representation of their flora and fauna, and to gather and register an immense mythical corpus that reveals the rich world-view of this small group. The ethnological work combined the most pertinent qualitative techniques from anthropology with scientific rigour in observations. It also stressed the importance of fieldwork based on the method of participant observation, in line with the works of the Polish anthropologist Bronis{\\l}aw Malinowski (Malinowski 1922).  The Tom\\'ar\\^aho (or Tomar\\'axo) form a small ethnic group, which together with the Ybyt\\'oso (or Ebid\\'oso), makes up a larger group, the Ishir, known in Paraguay as the Chamacoco. The Chamacoco, traditionally hunter-gatherers, are related linguistically to the Zamuco family. Another indigenous group, the Ayoreo, also belongs to the Zamuco linguistic family. We do not propose to discuss the etymology of these words and names: disparate proposals abound, based mainly upon the writings of European chroniclers, going back as far as {\\it Viaje al R{\\'\\i}o de la Plata, 1534--1554} by harquebusier Ulrico Schmidl (1510--1580). In different catalogues of languages and dialects, found within compilations of chronicles and other European texts, terms such as {\\it Timinabas}, later {\\it Timinaha}, and variants crop up, always in the Zamuco category. It is mentioned that the peoples who spoke these languages lived in the Chaco woods, inland, far from the Paraguay river. It is also said that these natives had not yet been `subdued by the Jesuit Mission'. We gather then that they are referring to the Tom\\'ar\\^aho, those natives whose descendants Sequera visited in the neighbouring area of San Carlos in 1986 and who, due to their ancient customs and their own idiosyncrasies, remained alienated from Paraguayan society for many years.  Shamanism is just as relevant among the Chamacoco as it is in other indigenous cultures. Vocal music is closely related to rituals carried out by both male and female shamans. These prominent members of the tribe, known as {\\it konsaha} or {\\it ahanak}, create their own repertoires based on dreams, called {\\it chyk\\^era}, which stimulate the creation of their poems, melodies and rhythms. The Chamacoco shamans try to dominate their dreams by turning them into a chant called {\\it teichu}. The production of these songs is highly personal. While they can be transmitted to other members of the community, they are generally performed in groups where several songs mix together and a very distinctive musical atmosphere is created, together with rattles ({\\it sonajas}) and whistles as instrumental accompaniment. Rattles called {\\it osecha} or {\\it pa{\\^\\i}k\\^ara} by the Chamacoco are made from gourds ({\\it calabazas} of the species {\\it Lagenaria siceraria}) or with tortoise shells called {\\it enermitak} ({\\it Geochelone carbonaria}). Dry seeds or pebbles are put inside so as to produce sound. These rattles represent the sky and the shamans identify their upper part with the centre of the sky, {\\it porr hos\\'ypyte}. Painted on the body of the instrument are `visual narratives' in the shape of rhomboids, and inside these geometrical figures are representations of the stars, {\\it porrebija} (Sequera 2006). Studies suggest a close relationship between the various Chamacoco vocal and instrumental techniques---especially among the Tom\\'ar\\^aho, where shamanism has lasted longer and more intensively---and their vision of the world, a bond between musical expression and the natives' view of the universe that still needs to be researched in detail (see also Cordeu 1994).  Chamacoco shamanic practices are very similar to those of other South American indigenous cultures, and even to those of other continents: visionary dreams, personality split, trance achieved through chants, etc. are situations to be repeated in time and space. It is through these dreams that the {\\it konsaha} discover for the rest of the members of the tribe the true `topography' of the indigenous universe and the way it interrelates with the mythical tales inherited from their ancestors.\\footnote{A detailed anthropological analysis of the case of the {\\it Qom}, or {\\it Tobas}, from the Argentinian Chaco can be consulted in Wright (2005).}  \\vspace*{-0.5 cm} ",
        "conclusions": "In this article we have given a brief sketch of the immense richness of Chamacoco conceptions of their relationship with nature and the sky. Painstaking ethnographic work has brought to light the most remarkable features of this aboriginal group as a whole, and particularly the characteristics of the Tom\\'ar\\^aho and Ybyt\\'oso. Through this work increasing numbers of people have become familiar with their culture and come to develop a great respect for their traditions and cultural identity. The transcription of their language, and their slow social improvement and education, are some of the elements that will help others to build upon this work in the future.  In the course of this vital ethnographic work, however, astronomical knowledge and practices have not been treated as thoroughly as the other aspects. Among the things that remain to be done are a detailed analysis of the recognition and identification of outstanding aspects of the Chamacoco sky, and an in-depth study of the true meaning of the {\\it ch\\^aro}, the cosmic tree that once united the sky and the earth and which remains, even now, central to rituals concerning the origin of the world. Other aspects needing greater attention are the names of, and the stories that were almost certainly told about, prominent stars visible at different times of the year, such as Sirius, and prominent asterisms such as Orion's belt; imagery concerning the presence and characteristics of the Milky Way; and their interpretation of certain exceptional and sudden phenomena, such as total solar eclipses, and of objects that make prolonged appearances in the sky, such as comets, or sporadic ones such as meteors.  The story of {\\it Iodhle} (Venus), who long ago married a young Tom\\'ar\\^aho, is just one of many stories with astronomical elements that old people used to tell. Many of these stories are still remembered by older members of this group. It is necessary and urgent to try to incorporate aspects of this culture into the intangible heritage of humanity before it is lost forever. Another project being planned is to try to understand the relationship between the vocal and instrumental techniques of the Chamacoco and their view of the universe surrounding them, especially in the use they make of rattles ({\\it pa{\\^\\i}k\\^ara}) which, as we have seen, represent the starry sky.  As part of our project we intend to work together with members of the present Tom\\'ar\\^aho community to gather data on the recognition of stars and constellations and their distribution, dark zones in the sky, nebulae etc. These data would be presented as annotated maps that would convey the symbolic significance of the Tom\\'ar\\^aho sky of the Upper Chaco region at different times of the year. This could tell us a great deal about aboriginal imagery and also about many aspects of their everyday life that they project into the dark depths of the sky. In short, we consider that the Tom\\'ar\\^aho conception of the sky has not been sufficiently explored as a study in itself, and that more fieldwork would be welcome in the area of ethnoastronomy as an interdisciplinary activity that includes both anthropologists and astronomers."
    },
    "1107/1107.2665_arXiv.txt": {
        "abstract": "Rapid detection of compact binary coalescence (CBC) with a network of advanced gravitational-wave detectors will offer a unique opportunity for multi-messenger astronomy.  Prompt detection alerts for the astronomical community might make it possible to observe the onset of electromagnetic emission from CBC.  We demonstrate a computationally practical filtering strategy that could produce early-warning triggers before gravitational radiation from the final merger has arrived at the detectors. ",
        "introduction": "As a compact binary system loses energy to gravitational waves (\\GW{}s), its orbital separation decays, leading to a runaway inspiral with the \\GW\\ amplitude and frequency increasing until the system eventually merges.  If a neutron star (\\NS) is involved, it might become tidally disrupted near the merger and fuel an electromagnetic (\\EM) counterpart \\citep{shibata:2007}. Effort from both the \\GW\\ and the broader astronomical communities might make it possible to use \\GW\\ observations as early warning triggers for \\EM\\ follow-up. In the first generation of ground-based laser interferometers, the \\GW\\ community initiated a project to send alerts when potential \\GW\\ transients were observed in order to trigger follow-up observations by \\EM\\ telescopes.  The typical latencies were 30 minutes \\citep{HugheyGWPAW2011}. This was an important achievement, but too late to catch any prompt optical flash or the onset of an on-axis optical afterglow.  Since the \\GW\\ signal is in principle detectable even before the tidal disruption, one might have the ambition of reporting \\GW\\ candidates not minutes after the merger, but seconds before.  We explore one essential ingredient of this problem, a computationally inexpensive latency free, real-time filtering algorithm for detecting inspiral signals in \\GW\\ data.  We also consider the prospects for advanced \\GW\\ detectors and discuss other areas of work that would be required for rapid analysis.  Compact binary coalescence (\\CBC) is a plausible progenitor for most short gamma-ray bursts \\citep[short \\GRB{}s;][]{Lee:2005, nakar07}, but the association is not iron-clad \\citep{2011ApJ...727..109V}. The tidally disrupted material falls onto the newly formed, rapidly spinning compact object and is accelerated in jets along the spin axis with a timescale of $0.1$--$1$~s after the merger \\citep{Janka1999}, matching the short \\GRB\\ duration distribution well. Prompt \\EM\\ emission including the \\GRB\\ can arise as fast outflowing matter collides with slower matter ejected earlier in inner shocks. The same inner shocks, or potentially reverse shocks, can produce an accompanying optical flash \\citep{Sari99}. The prompt emission is a probe into the extreme initial conditions of the outflow, in contrast with afterglows, which arise in the external shock with the local medium and are relatively insensitive to initial conditions. Optical flashes have been observed for a handful of long \\GRB{}s \\citep{2011CRPhy..12..255A} by telescopes with extremely rapid response or, in the case of GRB 080319b, by pure serendipity, where several telescopes were already observing the afterglow of another \\GRB\\ in the same field of view \\citep[FOV;][]{2008Natur.455..183R}. The observed optical flashes peaked within tens of seconds and decayed quickly. For short \\GRB\\ energy balance and plasma density, however, the reverse shock model predicts a peak flux in radio, approximately 20 minutes after the \\GRB, but also a relatively faint optical flash \\citep{nakar07}; for a once-per-year Advanced LIGO event at $130$~Mpc, the radio flux will peak around 9~GHz at $\\sim$$5$~mJy, with emission in the $R$-band at $\\sim$19~mag. Interestingly, roughly a quarter to half of the observed short \\GRB{}s also exhibit extended X-ray emission of $30$--$100$~s in duration beginning $\\sim$$10$~s after the \\GRB\\ and carrying comparable fluence to the initial outburst. This can be explained if the merger results in the formation of a proto-magnetar that interacts with ejecta~\\citep{Bucciantini2011}. Rapid \\GW\\ alerts would enable joint \\EM\\ and \\GW\\ observations to confirm the short \\GRB-\\CBC\\ link and allow the early \\EM\\ observation of exceptionally nearby and thus bright events.  In 2010 October, \\LIGO{}\\footnote{\\url{http://www.ligo.org/}} completed its sixth science run (S6) and Virgo\\footnote{\\url{http://www.ego-gw.it/}} completed its third science run (VSR3).  While both \\LIGO\\ detectors and Virgo were operating, several all-sky detection pipelines operated in a low-latency configuration to send astronomical alerts, namely, Coherent WaveBurst (cWB), Omega, and Multi-Band Template Analysis \\citep[MBTA;][]{HugheyGWPAW2011, S6lowlatency2, S6lowlatency3, S6lowlatency4}.  cWB and Omega are both unmodeled searches for bursts based on time-frequency decomposition of the GW data.  \\mbta\\ is a novel kind of template-based inspiral search that was purpose-built for low latency operation.  \\mbta\\ achieved the best \\GW\\ trigger-generation latencies of 2--5 minutes. Alerts were sent with latencies of 30--60 minutes, dominated by human vetting. Candidates were sent for \\EM\\ follow-up to several telescopes; \\textit{Swift}, LOFAR, ROTSE, TAROT, QUEST, SkyMapper, Liverpool Telescope, Pi of the Sky, Zadko, and Palomar Transient Factory \\citep{kanner2008, HugheyGWPAW2011} imaged some of the most likely sky locations.  There were a number of sources of latency associated with the search for \\CBC\\ signals in S6/VSR3 \\citep{HugheyGWPAW2011}, listed here.  \\paragraph{Data acquisition and aggregation ($\\gtrsim$100~ms)}% The \\LIGO\\ data acquisition system collects data from detector subsystems 16 times a second~\\citep{Bork2001}. Data are also copied from all of the \\GW\\ observatories to the analysis clusters over the Internet, which is capable of high bandwidth but only modest latency.  Together, these introduce a latency of $\\gtrsim$$100$~ms.  These technical sources of latency could be reduced with significant engineering and capital investments, but they are minor compared to any of the other sources of latency.  \\paragraph{Data conditioning ($\\sim$1~min)}% Science data must be calibrated using the detector's frequency response to gravitational radiation.  Currently, data are calibrated in blocks of 16~s. Within $\\sim$1~minute, data quality is assessed in order to create veto flags. These are both technical sources of latency that might be addressed with improved calibration and data quality software for advanced detectors.  \\paragraph{Trigger generation (2--5~min)}% Low-latency data analysis pipelines deployed in S6/VSR3 achieved an impressive latency of minutes.  However, second to the human vetting process, this dominated the latency of the entire \\EM\\ follow-up process.  Even if no other sources of latency existed, this trigger generation latency is too long to catch prompt or even extended emission. Low-latency trigger generation will become more challenging with advanced detectors because inspiral signals will stay in band up to 10 times longer.  In this work, we will focus on reducing this source of latency.  \\paragraph{Alert generation (2--3~min)}% S6/VSR3 saw the introduction of low-latency astronomical alerts, which required gathering event parameters and sky localization from the various online analyses, downselecting the events, and calculating telescope pointings.  If other sources of latency improve, the technical latency associated with this infrastructure could dominate, so work should be done to improve it.  \\paragraph{Human validation (10--20~min)}% Because the new alert system was commissioned during S6/VSR3, all alerts were subjected to quality control checks by human operators before they were disseminated. This was by far the largest source of latency during S6/VSR3. Hopefully, confidence in the system will grow to the point where no human intervention is necessary before alerts are sent, so we give it no further consideration here.  \\paragraph{}  This work will focus on reducing the latency of trigger production.  Data analysis strategies for advance detection of \\CBC{}s will have to strike a balance between latency and throughput. \\CBC\\ searches consist of banks of matched filters, or cross-correlations between the data stream and a bank of nominal ``template'' signals.  There are many different implementations of matched filters, but most have high throughput at the cost of high latency, or low latency at the cost of low throughput.  The former are epitomized by the overlap-save algorithm for frequency-domain (\\FD) convolution, currently the preferred method in \\GW\\ searches.  The most obvious example of the latter is direct time domain (\\TD) convolution, which is latency-free.  However, its cost in floating point operations per second is linear in the length of the templates, so it is prohibitively expensive for long templates.  The computational challenges of low-latency \\CBC\\ searches are still more daunting for advanced detectors for which the inspiral signal remains in band for a large fraction of an hour~(see the Appendix).  Fortunately, the morphology of inspiral signals can be exploited to offset some of the computational complexity of known low-latency algorithms.  First, the signals evolve slowly in frequency, so that they can be broken into contiguous band-limited time intervals and processed at possibly lower sample rates. Second, inspiral filter banks consist of highly similar templates, admitting methods such as the singular value decomposition (\\SVD)~\\citep{Cannon:2010p10398} or the Gram-Schmidt process~\\citep{rbf} to reduce the number of templates.  Several efforts that exploit one or both of these properties are under way to develop low-latency \\CBC\\ search pipelines with tractable computing requirements. One example is \\mbta{}~\\citep{Marion2004, Buskulic2010}, which was deployed in S6/VSR3.  MBTA consists of multiple, usually two, template banks for different frequency bands, one which is matched to the early inspiral and the other which is matched to the late inspiral.  An excursion in the output of any filter bank triggers coherent reconstruction of the full matched filtered output.  Final triggers are built from the reconstructed matched filter output.  Another novel approach using networks of parallel, second-order infinite impulse response~(IIR) filters is being explored by \\citet{shaunIIR} and \\citet{linqingIIR}.  We will use both properties to demonstrate that a very low latency detection statistic is possible with current computing resources.  Assuming the other technical sources of latency can be reduced significantly, this could make it possible to send prompt alerts to the astronomical community.  The paper is organized as follows.  First, we discuss prospects for early-warning detection.  Then, we provide an overview of our novel method for detecting CBC signals near real-time. We then describe a prototype implementation using open source signal processing software.  To validate our approach we present a case study focusing on a particular subset of the \\NS--\\NS\\ parameter space.  We conclude with some remarks on what remains to prepare for the advanced detector era. ",
        "conclusions": "We have demonstrated a computationally feasible filtering algorithm for the rapid and even early-warning detection of \\GW{}s emitted during the coalescence of \\NS{}s and stellar-mass black holes.  It is one part of a complicated analysis and observation strategy that will unfortunately have other sources of latency.  However, we hope that it will motivate further work to reduce such technical sources of \\GW\\ observation latency and encourage the possibility of even more rapid \\EM\\ follow-up observations to catch prompt emission in the advanced detector era.  \\CBC\\ events may be the progenitors of some short hard \\GRB{}s and are expected to be accompanied by a broad spectrum of \\EM\\ signals. Rapid alerts to the wider astronomical community will improve the chances of detecting an \\EM\\ counterpart in bands from gamma-rays down to radio. In the Advanced LIGO era, it appears possible to usefully localize a few rare events prior to the \\GRB{}, allowing multi-wavelength observations of prompt emission. More frequently, low-latency alerts will be released after merger but may still yield extended X-ray tails and early on-axis afterglows.  The \\lloid\\ method is as fast as conventional \\fft-based, \\FD\\ convolution but allows for latency free, real-time operation.  We anticipate requiring $\\gtrsim$40 modern multi-core computers to search for binary \\NS{}s using coincident \\GW\\ data from a four-detector network.  In the future, additional computational savings could be achieved by conditionally reconstructing the \\SNR\\ time series only during times when a composite detection statistic crosses a threshold~\\citep{svd-compdetstat}.  However, the anticipated required number of computers is well within the current computing capabilities of the \\LIGO\\ Data Grid.  We have shown a prototype implementation of the \\lloid\\ algorithm using \\gstreamer, an open-source signal processing platform.  Although our prototype already achieves latencies of less than one second, further fine tuning may reduce the latency even further.  Ultimately the best possible latency would be achieved by tighter integration between data acquisition and analysis with dedicated hardware and software. This could be considered for third-generation detector design.  Also possible for third-generation instruments, the \\lloid\\ method could provide the input for a dynamic tuning of detector response via the signal recycling mirror to match the frequency of maximum sensitivity to the instantaneous frequency of the \\GW\\ waveform.  This is a challenging technique, but it has the potential for substantial gains in \\SNR\\ and timing accuracy \\citep{PhysRevD.47.2184}.  Although we have demonstrated a computationally feasible statistic for advance detection, we have not yet explored data calibration and whitening, triggering, coincidence, and ranking of GW candidates in a framework that supports early \\EM\\ follow-up.  One might explore these and also using the time slice decomposition and the \\SVD\\ to form low-latency signal-based vetoes (e.g.,~$\\chi^2$~statistics) that have been essential for glitch rejection used in previous \\GW\\ \\CBC\\ searches.  These additional stages may incur some extra overhead, so computing requirements will likely be somewhat higher than our estimates.  Future work must more deeply address sky localization accuracy in a realistic setting as well as observing strategies. Here, we have followed \\citet{Fairhurst2009} in estimating the area of 90\\% localization confidence in terms of timing uncertainties alone, but it would be advantageous to use a galaxy catalog to inform the telescope tiling \\citep{galaxy-catalog}. Because early detections will arise from nearby sources, the galaxy catalog technique might be an important ingredient in reducing the fraction of sky that must be imaged.  Extensive simulation campaigns incorporating realistic binary merger rates and detector networks will be necessary in order to fully understand the prospects for early-warning detection, localization, and \\EM\\ follow-up using the techniques we have described."
    },
    "1107/1107.0660_arXiv.txt": {
        "abstract": "{We present here the results of a 180~ks \\Chan{}-LETGS observation as part of a large multi-wavelength campaign on Mrk 509.} {We study the warm absorber in Mrk 509 and use the data from a simultaneous \\COS{} observation in order to assess whether the gas responsible for the UV and X-ray absorption are the same.} {We analyzed the LETGS X-ray spectrum of Mrk 509 using the {{\\tt SPEX}} fitting package.} {We detect several absorption features originating in the ionized absorber of the source, along with resolved emission lines and radiative recombination continua. The absorption features belong to ions with, at least, three distinct ionization degrees. The lowest ionized component is slightly redshifted ($\\Delta v = 73$~km~s$^{-1}$) and is not in pressure equilibrium with the others, and therefore it is not likely part of the outflow, possibly belonging to the interstellar medium of the host galaxy. The other components are outflowing at velocities of $-196$ and $-455$~km s$^{-1}$, respectively. The source was observed simultaneously with \\COS{}, finding 13 UV kinematic components. At least three of them can be kinematically associated with the observed X-ray components. Based on the \\COS{} results and a previous {{\\it FUSE}} observation, we find evidence that the UV absorbing gas might be co-located with the X-ray absorbing gas and belong to the same structure.} {} ",
        "introduction": "\\label{intro}  Active galactic nuclei (AGN) are among the most luminous persistent extragalactic sources known in the Universe. A significant fraction of their bolometric luminosity is emitted at X-ray wavelengths, and originates in the innermost regions of these sources, where a supermassive black hole resides. Gravitational accretion of matter onto this black hole is believed to be the mechanism that powers these objects, and the intense radiation field thus generated might drive gas outflows from the nucleus. Indeed, more than 50\\% of the Seyfert 1 galaxies presents evidence of a photoionized gas, the so-called warm absorber (WA hereafter), in their soft X-ray spectra (e.g. Reynolds \\cite{Rey97}, George et al. \\cite{Geo98}, Crenshaw et al. \\cite{Cre03}). These WA are usually blueshifted with respect to the systemic velocity of the source with velocities of the order of hundreds of km~s$^{-1}$ (Kaastra et al. \\cite{Kaa00}, Kaspi et al. \\cite{Kas01}, Kaastra et al. \\cite{Kaa02}).  Interestingly, sources that present absorption in the UV also show absorption in X-rays (Crenshaw et al. \\cite{Cre99}), which may point to a possible relation between the two phenomena. With the advent of modern medium- and high-resolution spectrographs, the ionization state, kinematic properties and column densities of the absorption features in X-rays could be determined with higher accuracy. However, even though the average outflow velocity measured in the X-rays is generally consistent with the velocity shifts measured in the UV, the spectral resolution in X-rays is insufficient to resolve the many individual kinematic components present in the UV.  Some authors discussed the possibility that at least a fraction of the UV absorption is produced in the same gas responsible for the X-ray absorption (Mathur et al. \\cite{Mat94}, \\cite{Mat95}, \\cite{Mat99}; Crenshaw \\& Kraemer \\cite{CK99}, Kriss et al. \\cite{Kriss00}, Kraemer et al. \\cite{Kra02}), but the results were inconclusive and the connection between the two absorbing components seems to be very complex. The most common source of uncertainty came from the lack of correspondence between the column densities of the species detected in both UV and X-rays. However, more recent high-resolution observations indicate that these values can be reconciled if one assumes a velocity-dependent partial covering in the UV (e.g. Arav et al. \\cite{Arav02}, \\cite{Arav03}).  Simultaneous UV/X-ray observations are therefore the key to draw a global and unbiased picture of the WA in AGN (see e.g. Costantini \\cite{Cos10} for a review). High-resolution UV observations covering a wide range of ionization states are essential to determine whether some of the UV absorbers are indeed associated with the X-ray warm absorber. Low-ionization absorbers, such as \\MgII{} and \\CII{}, are only detected in the UV domain, while high-ionization ions like \\OVIII{} or \\MgXI{} can only be detected in X-rays. However, there is a range of ionization states that may produce ions detected in both bands, namely \\NV{} and, particularly, \\OVI{}. Therefore, several simultaneous or quasi-simultaneous X-ray and UV spectroscopic observations have been carried out on many sources known to host a X-ray WA and UV absorbers, e.g. NGC 3783 (Kaspi et al. \\cite{Kas02}, Gabel et al. \\cite{Gab03}), NGC 5548 (Crenshaw et al. \\cite{Cre03b}, Steenbrugge et al. \\cite{Ste05}), Mrk 279 (Arav et al. \\cite{Arav05}, Gabel et al. \\cite{Gab05}, Costantini et al. \\cite{Cos07}).  Mrk 509 is a Seyfert 1/QSO hybrid and is one of the best studied local AGN because of its high luminosity (L(1-1000~Ryd)$\\simeq 3.2\\times 10^{45}$~erg~s$^{-1}$) and proximity. It is particularly suitable for an extensive multi-wavelength campaign because of its brightness and confirmed presence of an intrinsic WA (Pounds et al. \\cite{Pou01}, Yaqoob et al. \\cite{Yaq03}, Smith et al. \\cite{Smi07}, Detmers et al. \\cite{Det10}) in X-rays, as well as a variety of UV absorption lines (Crenshaw et al. \\cite{Cre95}, \\cite{Cre99}; Savage et al. \\cite{Sav97}, Kriss et al. \\cite{Kriss00}, Kraemer et al. \\cite{Kra03}) and slow variability, which makes it particularly suited to reverberation studies.  This paper is devoted to the analysis of the \\Chan{}-LETGS observation of Mrk 509, which is part of an unprecedented multi-wavelength campaign of observations carried out at the end of 2009, aiming at determining the location of the warm absorber and the physics of its ionized outflow. The campaign involved five satellites (\\XMM{}, \\Chan{}, INTEGRAL, HST and Swift), and two ground-based facilities (WHT and PAIRITEL), thus covering more than five orders of magnitude in frequency. An overview of the campaign and its goals can be found in Kaastra et al. (\\cite{Kaa11a}), hereafter Paper I, and references therein.  This paper is organised as follows: in Sect.~\\ref{xrays} we briefly present the reduction method applied, in Sect.~\\ref{analysis} we describe the different models used to fit the spectrum, while in Sect.~\\ref{discussion} we discuss our results and compare them with previous X-ray observations as well as with the simultaneous \\COS{} observation. Finally, our conclusions are summarized in Sect.~\\ref{conclusions}. The quoted errors refer to 68.3\\% confidence level ($\\Delta \\chi^2 = 1$ for one parameter of interest) unless otherwise stated. ",
        "conclusions": "\\label{conclusions}  We have presented here the results from the 180~ks \\Chan{}-LETGS spectrum of Mrk 509 as part of an extensive multi-wavelength campaign. We found that the ionized absorber intrinsic to the source presents three distinct degrees of ionization ($\\log \\xi = $1.06, 2.26, and 3.15, respectively) and is somewhat shallow ($N_{\\rm H} = $ 1.9, 5.4, and 5.9 times $10^{20}$~cm$^{-2}$, respectively). The low-ionization component is slightly redshifted ($\\Delta v = 73$~km~s$^{-1}$), whereas the others are blueshifted with respect to the systemic velocity of the source ($\\Delta v = -196$ and $-$455~km~s$^{-1}$, respectively, see Table~\\ref{WApar}), and therefore they are outflowing. Their outflow velocities are correlated with the ionization parameter so that the higher the ionization parameter the faster is the outflow. The redshifted component is not in pressure equilibrium with the outflowing components, thus indicating that it may not form part of the same structure. We put mild constraints on the location of the outflowing gas, which would be consistent with a torus wind origin, although an inner origin cannot be ruled out. The energetics of the outflows indicate that they have a negligible effect in the interstellar medium of the host galaxy.  Mrk 509 was also observed in the UV domain with \\COS{} simultaneously with \\Chan{}, showing 13 kinematic components in the intrinsic UV absorption features. We compared the X-ray and UV absorbers in order to ascertain whether they have a common origin. At least three of the UV components could correspond with the X-ray components detected in the LETGS spectrum based only on the kinematics of the gas, which points to a possible co-location of both absorbing gases. The column density of \\OVI{} measured in one of the X-ray components is in agreement with that of a previous UV observation with {\\it FUSE} for their corresponding kinematic component. Similarly, the upper limits on the column densities of \\CIV{} and \\NV{} found in X-rays are also consistent with the values obtained by \\COS{}. Therefore, the similar kinematics and column densities indicate that the gases responsible for the X-ray and UV absorption share a common origin."
    },
    "1107/1107.2453_arXiv.txt": {
        "abstract": "We present a one-port calibration technique for characterization of beam waveguide components with a vector network analyzer. This technique involves using a set of known delays to separate the responses of the instrument and the device under test. We demonstrate this technique by measuring the reflected performance of a millimeter-wave variable-delay polarization modulator. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5045_arXiv.txt": {
        "abstract": "Using Constrained Local UniversE Simulations (CLUES) of the formation of the Local Group in a cosmological context we investigate the recently  highlighted problem that the majority of the most massive dark subhaloes of the Milky Way are too dense to host any of its bright satellites. In particular, we examine the influence of baryonic processes and find that they leave a twofold effect on the relation between the peak of the rotation curve and its position (\\Vmax\\ and \\Rmax). Satellites with a large baryon fraction experience adiabatic contraction thus decreasing \\Rmax\\ while leaving \\Vmax\\ more or less unchanged. Subhaloes with smaller baryon fractions undergo a decrease in \\Vmax\\ possibly due to outflows of material. Furthermore, the situation of finding subhaloes in simulations that lie outside the confidence interval for possible hosts of the bright MW dwarf spheroidals, appears to be far more prominent in cosmologies with a high $\\sigma_8$ normalisation and depends on the mass of the host. We conclude that the problem cannot be simply solved by including baryonic processes and hence demands further investigations. ",
        "introduction": "\\label{sec:introduction} The $\\Lambda$ Cold Dark Matter (\\LCDM) model, first explored more than two decades ago \\citep{Davis85}, has been very successful in explaining a multitude of observations at cosmological scales, such as anisotropies of Cosmic Microwave Background radiation (CMB) \\citet[e.g.][]{Jarosik11} and galaxy clustering on large scales \\citep[e.g.][]{Cole05}. However, on smaller, galactic scales the tests of the \\LCDM\\ model are complicated by the baryonic physics involved in galaxy formation. Therefore, testing the currently accepted concordance model at these scales is necessary in order to not only understand the nature of dark matter but also the accuracy of the model itself.  The validity of the \\LCDM\\ model on galactic scales is still being questioned due to the discrepancy between the number of observed satellites and  the number of predicted dark matter subhaloes. High resolution simulations of galactic-size haloes resolve a substantial number of substructures within the virial radius, as first pointed out by \\citet{Klypin99} and \\citet{Moore99}, and recently reviewed by \\citet{Kravtsov10} and \\citet{Bullock10}.  The most popular interpretation of this so-called \"Missing Satellite Problem\" requires that the smallest dark matter haloes are inefficient at forming stars \\citep[e.g.][]{Bullock10,Kravtsov10}. Mechanisms such as early reionization of the intergalactic medium and supernovae feedback have been invoked to identify the halo mass scale where the galaxy formation starts to be inefficient \\citep{Bullock00,Somerville02,Benson02}, partially solving the problem. Furthermore, the detection of satellites is most certainly biased because of current detection limits \\citep{Tollerud08,Walsh09}.  There is yet another aspect of the satellite population that needs to be addressed: the mismatch between the predicted and inferred distribution of \\Vmax\\ values at the high-\\Vmax\\ end as recently highlighted by \\citet{Boylan11}, where \\Vmax\\ measures the peak of the rotation curve of subhaloes. Using the Aquarius simulations \\citep{Springel08} and the Via Lactea II simulation \\citep{Diemand08} they found that the majority of the most massive subhaloes (i.e. the high-\\Vmax\\ objects) of the Milky Way are too dense to host any of its bright satellites.  There are a number of ways in which this discrepancy may be resolved: the subhalo mass function of the Milky Way could be a statistical anomaly with respect to the \\LCDM\\ expectations \\citep{Liu10,Guo11}, or the fundamental assumption that the luminosities of the satellites are not monotonically related to the mass of the subhaloes does not hold true.  In response to the claims by \\citet{Boylan11}, \\citet{Lovell11} explored the possibility that warm rather than cold dark matter can provide a better match to the inferred distribution of satellite circular velocities. With a power spectrum suppressed at masses below $\\sim10^{10}M_\\odot$ (corresponding to a warmon mass of 2~keV), they found that a warm dark matter model naturally produces haloes that are less concentrated than their cold dark matter counterparts. The attempt to explain the evolution of small scale structures in the local universe with a \\LWDM\\ model was already presented in \\citet{TikhonovGott09}. However, this is only one possible solution to the problem.  Baryonic processes will most certainly also affect the dark matter distribution. \\citet{Blumenthal86} showed that dissipative baryons will lead directly to the adiabatic contraction of the halo increasing its central density, thus being a critical ingredient to determine subhalo properties. However, the possibility that the influence of baryons will lead to a flattening of the dark matter central density cusp (through dynamical friction of infalling substructures composed of dark matter \\textit{and} baryons) has, for instance, been suggested by \\citet{El-Zant01} and further studied in \\citet{Romano-Diaz08}. Another way in which the haloes' density can be reduced is through sudden mass outflows that can alter substantially the central structure, as suggested by \\citet{Navarro96b}. In a recent work of \\citet{Parry11} this last scenario has been tested by following the evolution of one simulated satellite, with promising results. The same authors though also showed that the inclusion of baryons in simulations does not seem to have any correlation with the increase or decrease of the dark matter central density.  In this \\textit{Letter} we directly address the issue of the \\Vmax\\ problem in \\LCDM\\ simulations by comparing two identical simulations with each other: one that is solely based upon dark matter physics and another incorporating all the relevant baryonic physics. These simulations form part of the CLUES project\\footnote{\\texttt{http://www.clues-project.org}}, in which the initial conditions are set by imposing constraints derived from observational data of the Local Group. The main feature of using constrained simulations is that it provides a numerical environment that closely matches our actual neighborhood. ",
        "conclusions": "\\label{sec:conclusion}  In this letter we explored the possibility that baryonic processes may solve the recently presented problem of \"the puzzling darkness of Milky Way subhaloes\" \\citep{Boylan11}. To this extent, we used dark matter only as well as full hydrodynamical simulations of cosmic structure in the context of the CLUES project. We used cosmological parameters determined from both the WMAP3 and WMAP5 data.  Our conclusions are twofold and can be summarized as follows:  \\begin{itemize}   \\item We find that when baryonic physics is included, following the feedback and star formation prescriptions of \\cite{Springel03}, the problem of having too dense massive subhaloes is not solved. Instead, gasdynamical simulations pose new questions regarding which mechanisms are responsible for the lowering of \\Rmax\\ in those subhaloes (while \\Vmax\\ remains more or less constant). Adiabatic contraction seems to be a reasonable explanation, as shown using the modified adiabatic contraction model of \\citet{Gnedin04}: this process is effective only for some subhaloes, specifically, for those with a high baryon fraction. For the SPH subhaloes with lower baryon fractions at redshift $z=0$, instead, we observe a general increase of \\Rmax\\ with respect to their DM counterpart, thus meaning that other effects are at work, e.g. the model proposed by \\citet{Navarro96b} in which a rapid expulsion of baryonic mass during star formation causes a reduction of the halo concentration, as well as naturally explaining the low baryon fraction of these objects.   \\item While in the WMAP5 DM only case we find dark matter subhaloes outside the confidence area (calculated following the prescription given in \\citet{Boylan11}) in the WMAP3 cosmology we only have one massive subhalo outside this observational range. Since the Via Lactea II and Aquarius simulations presented in \\citet{Boylan11} are similar cases, we conclude that the cosmology certainly has an influence, too: the higher $\\sigma_8$ of the WMAP5 scenario eventually led to higher host masses which -- according to our test -- are the most likely reason for the higher number of excessively centrally concentrated substructures. Note that the latest data from WMAP7 favours $\\sigma_8=0.807$, a value between the WMAP3 and WMAP5 results: this could mean that the problem is worse than in WMAP3, but not as pronounced as in the WMAP5 case. \\end{itemize}  An issue neither touched upon by us nor other authors is the adequacy of using NFW profiles when calculating the confidence interval for possible hosts of the bright MW dwarf spheroidals. It is obvious that tidal effects will lead to severe modifications of the original NFW density profile subhaloes had upon infall into their host \\citep{Kazantzidis04}. They therefore leave an impact upon internal and kinematical properties, respectively \\citep{Lokas10,Lokas11}, which should be taken into account when using observed half-light radii $R_{1/2}$ and their corresponding masses $M_{1/2}$ to define the confidence interval. Further, \\citet{Romano-Diaz08} showed that adiabatic contraction makes the dark matter profile almost isothermal. However, the relevance is questionable as material will primarily be stripped from the outer regions: \\citet{Penarrubia08} state that dSphs embedded in NFW haloes are very resilient to tidal effects until they are nearly destroyed. This is supported by \\citet{Navarro10} who found that the NFW shape holds reasonably well even for subhaloes. To roughly gauge the problem, we fitted our (SPH) subhaloes to a NFW profile, and observed that while some of them are well fitted, there are still objects whose density profile cannot be approximated by the simple NFW functional form. Taking all these considerations into account suggests that the NFW profile used to calculate the allowed region is likely not the best choice.  The interpretation of the results presented here clearly demands a closer investigation of the evolutionary tracks of the satellites, the actual influence of the SF and feedback model as well as an improved calculation of the observational confidence level, verifying the applicability of the NFW approach. However, we leave this to a companion paper (Di Cintio et al., in preparation) and only highlight here that simply the inclusion of baryonic physics does not solve the problem; it rather poses new challenges to be explored and studied in greater detail."
    },
    "1107/1107.0726_arXiv.txt": {
        "abstract": "This paper is aimed at developing a better understanding of the structure of the information that in contained in galaxy surveys, so as  to find optimal ways to combine observables from such surveys. We first show how Jaynes' Maximal Entropy Principle allows us, in the general case, to express the Fisher information content of data sets in terms of the curvature of the Shannon entropy surface with respect to the relevant observables. This allows us to understand the Fisher information content of a data set, once a physical model is specified, independently of the specific way that the data will be processed, and without any assumptions of Gaussianity. This includes as a special case the standard Fisher matrix prescriptions for Gaussian variables widely used in the cosmological community, for instance for power spectra extraction. As an application of this approach, we evaluate the prospects of a joint analysis of weak lensing tracers up to second order in the shapes distortions, in the case that the noise in each probe can be effectively treated as model independent. These include the magnification, the two ellipticity and the four flexion fields. At the two point level, we show that the only effect of treating these observables in combination is a simple scale dependent decrease of the noise contaminating the accessible spectrum of the lensing E-mode. We provide simple bounds to its extraction by a combination of such probes, as well as its quantitative evaluation when the correlations between the noise variables for any two such probes can be neglected. ",
        "introduction": "With cosmological data sets currently going through a rapid period of growth, it is increasingly important to quantitatively understand the potential and limits of particular data sets to test a physical model or hypothesis. For this, the Fisher information matrix has become a widely used tool in cosmology.  The concept of Fisher information has a long history. It was first coined by the statistician and geneticist R.A. Fisher \\citep{fisher25} under the name of intrinsic accuracy of frequency curves. It has found its way into the cosmological community over the last decade, where it is often used to optimise survey configurations \\citep{2007MNRAS.381.1018A,2007MNRAS.377..185P,2006astro.ph..9591A,2009ApJ...695..652B} of planned cosmology experiments or to evaluate the expected errors on certain cosmological parameters with some observables \\citep{1997ApJ...480...22T,Tegmark97,1999ApJ...514L..65H,2004PhRvD..70d3009H}.  Much of the work to date has been limited to particular sets of observables and estimators. Usually, it is assumed that observational errors as well as the parameters probability distribution have Gaussian shape. The first aim of this work is to propose a framework to express the global Fisher information content of large data sets in a way that is independent of the specific ways that the data will be processed, and in realistic situations, where the exact statistical properties of the data are not known precisely. This should then provide a well motivated basis point in order to perform systematic and robust trade-off studies. For this purpose a number of useful concepts already exist in  the fields of information theory and probability theory, such as Shannon entropy or relative entropy \\citep{Kullback}, which we can use to gain a better understanding of what we can achieve with planned experiments. Specifically, we will show that we can achieve our aim by combining Fisher's information measure with Jaynes' Principle of Maximal Entropy.  \\citep{jaynes83,jaynes2003}. \\newline\\indent In a second step, as a concrete application of this approach, we investigate the joint entropy and information content of multiple observables of the same underlying, cosmologically interesting field. This is a very relevant situation in weak lensing \\citep{1992grle.book.....S,2001PhR...340..291B,2003ARA&A..41..645R,2006astro.ph.12667M,2006glsw.book.....S}, where the distortions of galaxy images to any order are sourced by the lensing potential field. \\newline\\newline\\indent This paper is divided into the following sections. In section \\ref{Setting the scene},  we present in details our approach.  We first review and develop some key properties of Fisher information, and its link to the Cramer Rao inequality. We put a strong emphasis on its interpretation as a measure of information on model parameters in a data set, that is obtained from the probability distribution of different observational outcomes as function of these same model parameters. Readers familiar with these aspects may jump to section {\\ref{The structure of the information in a physical model}}, where we introduce Jaynes' Maximum Entropy Principle, and show how it ideally completes Fisher's information measure, allowing us to understand the information content of a data set on a physical model in the case of incomplete knowledge. In section \\ref{application}, we show how the study of the Shannon entropy of a set of homogeneous fields provide a simple and model parameter independent answer to the question of the combination of the weak lensing observables shear, magnification and flexion. We provide in section \\ref{results}  quantitative evaluation of the prospects of such a combination at the two-point level for typical dark energy surveys parameters, and conclude in section \\ref{discussion} with a summary of the results and a discussion. A set of appendices collects some technical details. ",
        "conclusions": "\\label{discussion} We have shown how Jaynes'  Maximal Entropy Principle allows us to construct the Fisher information content on model parameters in a given data set in the form of equation (\\ref{sumform}) or (\\ref{fishermaxent}). This is done by making the key quantity the entropy of the distribution as function of the constraints that we put on it. These constraints form our knowledge of the statistical properties of the future data. To the best of the authors knowledge,  equation (\\ref{sumform}) or (\\ref{fishermaxent}) are not to be found in this form in the literature. However, they cannot  be considered new, since as stated earlier, they can be easily gained from the Fisher information content of the exponential family of distributions \\citep{Jennrich75,vandenbos07}, after the identification of the curvature of the entropy surface with the generalised inverse of the covariance matrix. Especially, the maximal entropy distributions are precisely those for which the Cram\\'{e}r-Rao inequality is an equality, since the curvature of the entropy surface is the inverse correlation matrix between the model predictions. Equation (\\ref{fishermaxent}) also bears a strong formal similarity to the well known result  \\cite[chap. 2]{Kullback} or \\cite[chap. 6]{caticha08}, that the Fisher information can always be written as the curvature of the Kullback-Leibler divergence for distributions parametrised by the same set of parameters. \\newline\\newline The Fisher matrices currently used in weak lensing or clustering can all be seen as special cases of this approach, namely equation (\\ref{FPk}), when knowledge of the statistical properties of the future data does not go beyond the two-point statistics. Indeed, in the case that the model does not predict the means, and knowing that for discrete fields the spectral matrices, equation (\\ref{spectralmatrix}), carry a noise term due to the finite number of galaxies, or, in the case of weak lensing, also due to the intrinsic ellipticities of galaxies, the amount of information in (\\ref{FPk})  is essentially identical to the standard expressions used to predict the accuracy with which parameters will be extracted from power-spectra analysis. \\newline\\newline There is, however, a conceptual difference worth noting in that the standard approach is to pick an estimator for the power-spectra and assume that both the fields as well as the distribution of the estimators are Gaussian. The result is the amount of Fisher information there is in the power spectra, under the assumptions of Gaussian statistics for the estimators and the fields. In our approach, the only assumption is on the fields distribution. Our results do not depend on the way the information will be extracted, but shows the amount of Fisher information in the fields as a whole. \\newline\\newline Of course, the maximal entropy approach, which tries to capture the relevant properties of $p_X$ through a sophisticated guess, gives no guaranties that its predictions are actually correct. Nevertheless, as discussed in section \\ref{redundant}, it provides a systematic approach with which to update the probability density function in case of improved knowledge of the relevant physics. \\newline \\newline Using this formalism we have investigated the combined Fisher information content of weak lensing probes up to second order in the shapes distortions, assuming model parameter independent noise. By having a look at the joint Shannon entropy of the fields, we have shown how the only effect of treating these observables jointly is to reduce the effective level of noise contaminating the convergence field, according to equations (\\ref{noise}) and (\\ref{Entropy2}), independently of the model parameters. \\newline\\newline\\indent These are the key points of this paper that we would like to emphasize : \\begin{enumerate} \\item Equation (\\ref{fishermaxent}) presents a measure of information content that depends only on the constraints put on the data and the physical model. It is written in terms of the curvature of Shannon's entropy surface for maximal entropy distributions. It can always be interpreted, regardless of the actual distribution of the parameters, and of the specific way the analysis will proceed, as the expected curvature of the $\\chi^2$ surface to the full set of model predictions. Assumptions of Gaussianity are neither needed nor used at any point. \\item  Over a very wide range of scales, the probe of choice both for mass reconstruction or cosmological purposes are the ellipticity components of the galaxies. Flexion takes over only on the arcsecond scale. In combination with the ellipticities, it can lead to substantial increase in statistical power on the scale of the arcminute. From the cosmological point of view, we expect size information to contribute at the $10\\%$ level of the total information content. The only key parameter is the combination (\\ref{improve}) of the dispersion values and the permeability $\\alpha_s$ of the population sizes to the convergence field. On the other hand, the prospects of including flexion in cosmological analysis are less clear. The most obvious drawback being the need for an accurate understanding of the non-linear power spectrum. \\item Besides, our results render the inclusion of flexion and size information within more detailed Fisher matrix analysis for future dark energy experiments extremely simple, such as in the exhaustive approach combining the information galaxy density fields with shear fields in a tomographic setting of \\citep{2009ApJ...695..652B}. From (\\ref{multis}) follows namely that the inclusion of all the two-point correlations of these additional weak lensing probes can be accounted for by adapting the noise term $N_\\eff$. \\end{enumerate} \\black Possible developments of this work includes the relaxation of the main limitation of the results, for instance the assumption that the noise is independent of the model parameters. Also, we plan to show that the approach presented in the first part of this work can lead to quantitative evaluations in non Gaussian cases as well, when other observables than the first two moments are considered."
    },
    "1107/1107.5932_arXiv.txt": {
        "abstract": "We report Hubble Space Telescope observations of 6 gravitational lenses with the Advanced Camera for Surveys. We measured the flux ratios between the lensed images in 6 filters from 8140\\AA\\ to 2200\\AA. In 3 of the systems, HE0512$-$3329, B1600+434, and H1413$+$117, we were able to construct UV extinction curves partially overlapping the 2175\\AA\\ feature and characterize the properties of the dust relative to the Galaxy and the Magellanic Clouds.  In HE1104$-$1804 we  detect chromatic microlensing and use it to study the physical properties of the quasar accretion disk. For a Gaussian model of the disk  $\\exp(-r^2/2 r_s^2)$, scaling with wavelength as $r_s \\propto \\lambda^p$, we estimate $r_s( \\lambda3363)=4^{+4}_{-2}$ ($7\\pm 4$)  light-days and  $p=1.1\\pm 0.6$ ($1.0\\pm 0.6$) for a logarithmic (linear) prior on $r_s$. The remaining two systems, FBQ0951+2635 and SBS1520+530, yielded no useful estimates of extinction or chromatic microlensing. ",
        "introduction": "The wavelength-dependent flux ratios between gravitationally lensed images of quasars can be used to probe both the extinction law in the lens galaxy and the unresolved structure of the quasar.  Each image is differentially extincted by dust along the  path of the image through the lens galaxy, and this can be used to measure the extinction, the extinction law, and the lens redshift (e.g. Nadeau et al. 1991; Falco et al. 1999; Motta et al. 2002; Mu\\~noz et al. 2004; Mediavilla et al. 2005; Eliasdottir et al. 2006, Mosquera et al. 2011). The second effect, microlensing by the stars in the lens galaxy (Chang \\& Refsdal 1979), leads to wavelength dependent changes in the flux ratios because the effective size of the quasar accretion disk varies with wavelength (Wisotzki et al. 1993, 1995; Wucknitz et al. 2003; Anguita et al. 2008; Bate et al. 2008; Eigenbrod et al. 2008; Poindexter et al.2008; Floyd et al. 2009; Mosquera et al. 2009, 2011; Blackburne et al. 2010; Mediavilla et al. 2011).  Detections of ``chromaticity'' between the images of a lensed quasar are useful for studying both phenomena if they can be disentangled.  In extragalactic astronomy, understanding dust is crucial to understanding galaxies, through its effects on estimates of star formation rates and galaxy evolution (e.g. Conroy et al. 2009), cosmology, through its effects on SNe~Ia fluxes (e.g. Jha et al. 2006), and the interpretation of gamma-ray burst afterglows (e.g. Jakobsson et al. 2004)).  Unfortunately, classical methods for obtaining accurate extinction curves to characterize dust cannot be used outside the Local Group because they depend on detailed measurements of individual stars.  Gravitational lenses are one of the best quantitative astrophysical probes of dust properties at intermediate redshifts given lenses with the right amount of dust and the appropriate combinations of redshifts. If there is too little dust, it is difficult to measure the extinction at long wavelengths and microlensing is more likely to dominate the chromaticity. If there is too much dust, it becomes impossible to measure extinction curves into the rest-frame ultraviolet.  Similarly, the lens redshift must be high enough to make the rest-frame 2175\\AA\\ dust feature observable, while the source redshift must be low enough to avoid having the quasar continuum blocked by absorption in the intergalactic medium.  Similar considerations hold for studying chromatic microlensing over the broadest possible wavelength baseline.  We selected six lenses from the survey of extinction by Falco et al. (1999) that roughly satisfied these criteria: HE~0512$-$3329, FBQ~0951$+$2635, HE~1104$-$1805, H~1413$+$117, SBS~1520$+$530, and B~1600$+$434. As we report in \\S2, we observed them in 6 filters spanning 2200\\AA\\ to 8100\\AA\\ (F220W, F250W, F330W, F435W, F555W, F625W and F814W) using  the Hubble Space Telescope (HST) and the ACS/HRC camera. This approach ensures that we can measure the image flux ratios without contamination from the lens or the host galaxy of the quasar.  Section 2 also outlines how we model the results to study extinction and microlensing. In \\S3 we present the results, reporting on the extinction curves of three of the systems and the chromatic microlensing in one system. Section 4 summarizes the results and lessons for future observations. ",
        "conclusions": "The effects of both extinction and microlensing become larger as we observe them at shorter wavelengths.   Unfortunately, the atmosphere prevents us from observing into the ultraviolet from the ground, and so we generally cannot observe the rest-frame 2175\\AA\\ region to search for the characteristic feature of Galactic and LMC extinction curves or to probe the hot regions near the inner edges of accretion disks. Here we surveyed six gravitational lenses with evidence for significant wavelength dependent flux ratios in the extinction study of Falco et al. (1999) from the I-band into the UV (8100\\AA\\ to 2200\\AA). Two of the lenses, FBQ~0951+2635 and SBS~1520+530, showed changes with wavelength that were too small to yield interesting constraints.  It is not surprising given the selection method that three of the lenses show significant evidence for differential extinction between the images. We argue for extinction dominating over chromatic microlensing systems based on the lack of evidence for significant time variability in the color, although all three systems show small changes in the flux ratios that are probably due to microlensing.  In the case of HE~0512$-$3329 we find evidence for a weak 2175\\AA\\ feature from the dust in the $z_l=0.93$ lens. For B~1600+434 we cannot quite reach the wavelengths needed to quantify the presence of the 2175\\AA\\ feature, although a CCM extinction law agrees with our observations, while the lack of a lens redshift for H~1413+117 limits our conclusions.  Both systems contain significant differential extinction, and it is likely that the dust in B~1600+434 has the 2175\\AA\\  feature and that the dust in H~1413+177 does not.  We clearly detect chromatic microlensing in HE~1104$-$1805.  If we estimate the wavelength dependent size of the accretion disk by modeling our single epoch of data or the earlier CASTLES data, we find  compatible results. If we combine the two single epoch estimates, the combined result agrees with the multi-band light curve analyses of Poindexter et al. (2008). Modeled as a Gaussian source $\\exp(-r^2/2 r_s^2)$ with $r_s \\propto \\lambda^p$ and normalized at the observed wavelength $\\lambda=3363$ \\AA\\  we find $r_s=4^{+4}_{-2}$ ($7\\pm 4$)  and $p=1.1\\pm 0.6$ ($1.0\\pm 0.6$) for a logarithmic (linear) prior on $r_s$. These slopes are consistent with the expected slopes from standard thin disk theory ($T\\propto R^{-1/p}$ with $p$=4/3), but the uncertainties are too large to draw a stronger conclusion.    \\bigskip"
    },
    "1107/1107.1516_arXiv.txt": {
        "abstract": "We present multi-epoch radio and optical observations of the M7 dwarf 2MASS J13142039+1320011.  We detect a $\\sim 1$ mJy source at 1.43, 4.86, 8.46 and 22.5 GHz, making it the most luminous radio emission over the widest frequency range detected from an ultracool dwarf to date.  A 10 hr VLA observation reveals that the radio emission varies sinusoidally with a period of $3.89\\pm 0.05$ hr, and an amplitude of $\\approx 30\\%$ at 4.86 GHz and $\\approx 20\\%$ at 8.46 GHz.  The periodicity is also seen in circular polarization, where at 4.86 GHz the polarization reverses helicity from left- to right-handed in phase with the total intensity.  An archival detection in the FIRST survey indicates that the radio emission has been stable for at least a decade.  We also detect periodic photometric variability in several optical filters with a period of $3.79$ hr, and measure a rotation velocity of $v{\\rm sin}i=45\\pm 5$ km s$^{-1}$, in good agreement with the radio and optical periods.  The period and rotation velocity allow us to place a lower limit on the radius of the source of $\\gtrsim 0.12$ R$_{\\odot}$, about $30\\%$ larger than theoretical expectations.  The properties of the radio emission can be explained with a simple model of a magnetic dipole mis-aligned relative to the stellar rotation axis, with the sinusoidal variations and helicity reversal due to the rotation of the magnetic poles relative to our line of sight.  The long-term stability of the radio emission indicates that the magnetic field (and hence the dynamo) is stable on a much longer timescale than the convective turn-over time of $\\sim 0.2$ yr.  If the radio emission is due to the electron cyclotron maser process, the inferred magnetic field strength reaches at least 8 kG. ",
        "introduction": "\\label{sec:intro}  An unexpected result from radio observations of low mass stars and brown dwarfs was the recent discovery of periodic radio emission tied to the stellar rotation \\citep{brr+05,had+06,hbl+07,brp+09}. Previously thought to be incapable of generating and dissipating strong and stable magnetic fields, observations in recent years indicate that at least 10\\% of fully convective late-M and L dwarfs possess stable and large-scale magnetic fields \\citep{ber02,ber06,bbf+10}.  When combined with X-ray and H$\\alpha$ observations \\citep{mb03,rb08,bbf+10}, Zeeman broadening of FeH molecular lines in Stokes $I$ \\citep{rb07}, and spectropolarimetry in Stokes $V$ (Zeeman Doppler Imaging: ZDI; \\citealt{dfc+06}), a picture of the magnetic topology of fully convective ultracool dwarfs is now beginning to emerge.  Zeeman broadening observations provide a measure the average surface magnetic field, $Bf$ (where $B$ is the magnetic field and $f$ is the covering fraction) and have been used to study moderately rotating objects as cool as M9 \\citep{rb10}. These observations uncovered average magnetic field strengths of up to a few kG.  The ZDI technique has been used to study the large scale structure and topology of magnetic fields (as it is sensitive only to the net polarization in Stokes $V$) in objects down to spectral type M8.  At the fully convective boundary (spectral type of $\\sim {\\rm M3}$) these studies hint at a transition from non-axisymmetric toroidal fields to predominantly poloidal, axisymmetric field topologies \\citep{dmp+08,mdp+08}, although the systematic uncertainties and general detectability of field topologies across a wide range in stellar parameters are not fully understood.  Similarly, there seems to be an increase in the ratio of magnetic fluxes measured in circularly polarized to unpolarized spectral lines at this boundary, apparently confirming that the fully convective objects have increasingly large-scale, low-multipole field topologies.  However, a number of fully convective mid-M dwarfs, potentially belonging to a younger population, still appear to possess small-scale complex fields \\citep{mdp+10}, although recent observations seem to question this idea \\citep{rei10}.  In the context of magnetic activity, the correlation between X-ray and radio emission, which holds across several orders of magnitude in each quantity and over a wide spectral type range \\citep{gb93,bg94}, sharply breaks down around spectral type M7--M8 \\citep{bbb+01,ber02,bbf+10}.  X-ray and H$\\alpha$ activity also decline precipitously around the same spectral type \\citep{mb03,rb08,bbf+10}.  Finally, the H$\\alpha$ and X-ray rotation-activity relations also break down near the M/L transition \\citep{mb03,bbg+08,rb10}.  The origin of these changes remains unclear, in particular whether they are related to a change in the magnetic dynamo itself, or to a change in the way magnetic energy is released in the cool and largely neutral atmospheres.  Radio activity in ultracool dwarfs, on the other hand, appears to be equally strong (or stronger when normalized to the bolometric luminosity) than in early-M dwarfs \\citep{ber02,ber06,bbf+10}.  Among the active population of ultracool dwarfs, the handful of objects with periodic radio emission exhibit diverse behavior.  In particular, the wide variation in the temporal behavior, emission bandwidth, and fractional polarization suggest that there are a number of different processes at work.  Three objects (TVLM\\,513-46546: \\citealt{had+06}; LSR\\,1835+3259: \\citealt{hbl+07}; and 2M\\,0746+2000: \\citealt{brp+09}) have been observed to emit short-duration, highly polarized pulses with apparently narrow bandwidth.  On the other hand, 2M\\,0036+1821 exhibits periodic variability with a different structure, rising and falling in a roughly sinusoidal manner \\citep{brr+05}. In all four objects, the observed periods, which range from two to three hours, are in close agreement with the known rotation velocities ($v{\\rm sin} i$).  Two primary emission mechanisms have been proposed to explain the observed radio emission: gyrosynchrotron radiation and the electron cyclotron maser (ECM) instability \\citep{ber06,tre06,had+08,brp+09}. The short-duration periodic pulses have been attributed to a coherent process such as ECM due their small bandwidth and high degree of circular polarization \\citep{hbl+07,brp+09}.  However, the mechanism for the sinusoidal radio emission from 2M\\,0036+1821, as well as the quiescent emission from several ultracool dwarfs, is still under debate, partly because of the low levels of circular polarization and the much wider bandwidth.  Here we present multi-wavelength observations of a newly discovered radio-emitting M7 dwarf, 2MASS J13142039+1320011 (hereafter, \\2m1314), which exhibits unique emission properties among the population of radio active ultracool dwarfs.  It has the brightest radio emission over the widest frequency range of any ultracool dwarf to date (detected from 1.4 to 22.5 GHz with a nearly flat spectrum), and presents the first clear example of sinusoidal emission (in phase at 4.96 and 8.46 GHz).  The periodicity is also detected in circular polarization, including a helicity reversal that provides insight into the topology of the emission regions.  Photometric observations reveal periodic variability with the same period as in the radio band, indicative of a large-scale magnetic spot structure.  The unique radio properties of \\2m1314, and its location at the cusp of the transition in magnetic field properties (spectral type M7), make it an invaluable laboratory for fully convective dynamos. ",
        "conclusions": "\\label{sec:conc}  \\2m1314\\ provides a unique window into magnetically-induced radio emission in ultracool dwarfs, coinciding with the spectral type at which the radio/X-ray correlation breaks down \\citep{ber02,ber06}.  It has the most luminous stable radio emission over the widest frequency range ($1.4-22.5$ GHz) to date, and furthermore exhibits clear sinusoidal periodicity, with the first example of a helicity reversal in the circularly polarized emission of an ultracool dwarf.  Taken in conjunction, these properties point to emission from a simple dipolar field, mis-aligned relative to the stellar rotation axis.  In the context of electron cyclotron maser emission, the magnetic field strength is required to be at least 8 kG.  The inferred field structure is similar to the simple field topologies inferred from Zeeman Doppler imaging of some fully convective dwarfs \\citep{mdp+10}, while the field strength in the case of ECM emission is larger than the typical average fields measured from Zeeman broadening of FeH lines \\citep{rb07}.  The stability of the radio emission for at least a decade indicates that the underlying dynamo is stable well beyond the convective turnover timescale for an M7 dwarf ($\\sim 70$ d; \\citealt{saa01}). Moreover, a comparison to convective dynamo models and simulations supports predictions of dominant large-scale, low-multipole fields \\citep{dsb06,bro08}, particularly for fast rotators \\citep{bro08}, but is at odds with models that predict small-scale, high-multipole fields \\citep{ddr93,ck06}.  Future observations of \\2m1314\\ will address several key questions. First, if the emission is due to the ECM process, we expect a cut-off in the radio emission at a frequency that corresponds to the maximal field strength, $\\nu_{\\rm max}=2.8\\times 10^6\\,B_{\\rm max}$ Hz. Measurements of this possible cut-off frequency will therefore provide a crucial constraint on the magnetic field since the Zeeman techniques are only sensitive to the average surface field or its large-scale component.  The increased sensitivity and bandwidth of the EVLA, particularly at frequencies of $\\sim 20-40$ GHz, is well-suited for this test.  Second, simultaneous long-term monitoring of the radio and optical emission will test whether the variability has a similar period and phase.  If identical, this will point to an origin in a related magnetic field structure, while a clear difference in period will be suggestive of differential rotation and an origin at different latitudes.  Finally, long-term astrometric monitoring with the VLBA will be able to uncover a companion down to a few Jupiter masses \\citep{fb09}.  Thus, further studies of \\2m1314\\ are bound to provide new insights into the properties of fully convective dynamos and the release of magnetic energy in the coolest stars and brown dwarfs."
    },
    "1107/1107.3808_arXiv.txt": {
        "abstract": "{The bulk of solar flare emission originates from very compact sources located in the lower solar atmosphere and seen in various wavelength ranges: near optical, UV, EUV,  soft and hard X-rays, and gamma-ray emission, yet very few spatially resolved imaging observations to determine the structure of these compact regions exist.} {We investigate the above-the-photosphere heights of hard X-ray (HXR), EUV and white-light  (6173 $\\AA$) continuum sources in the low atmosphere and the corresponding densities at these heights. Considering collisional transport of solar energetic electrons we also determine where and how much energy is deposited and compare these values with the emissions observed in HXR, EUV and continuum.} {Simultaneous EUV/continuum images from AIA/HMI on-board SDO and HXR RHESSI images are compared to study a well observed gamma-ray limb flare. Using RHESSI X-ray visibilities we determine the height of the HXR sources as a function of energy above the photosphere. Co-aligning  AIA/SDO and HMI/SDO images with RHESSI we infer, for the first time, the heights and characteristic densities of HXR, EUV and continuum (white-light) sources in the flaring footpoint of the loop.} {35-100 keV HXR sources are found at heights between 1.7 and 0.8 Mm above the photosphere, below the 6173 $\\AA$ continuum emission which appears at heights $1.5-3$ Mm, and the peak of EUV emission originating near 3 Mm. } {The EUV emission locations are consistent with energy deposition from low energy electrons of $\\sim 12$~keV occurring in the top layers of the fully ionized chromosphere/low corona and not by $\\gtrsim 20$ keV electrons that produce HXR footpoints in the lower neutral chromosphere. The maximum of white-light continuum emission appears between the HXR and EUV emission, presumably in the transition between ionized and neutral atmospheres suggesting free-bound and free-free continuum emission. We note that the energy deposited by low energy electrons is sufficient to explain the energetics of optical and UV emissions.} ",
        "introduction": "\\label{Introduction} Solar flare emission is now observed virtually over the whole electromagnetic spectrum from radio frequencies as low as $10-100$~MHz up to photon energies as high as a few hundred MeV. The bulk of this emission originates from the chromosphere, a small part of the solar atmosphere which is narrow in height and much smaller than the size of a flaring region. Being only about $<3$ Mm ($<4''$) across, the chromosphere presents a formidable observational challenge in both angular and temporal resolution especially for transient phenomena like solar flares. During solar flares large numbers of energetic electrons are accelerated in the solar corona and effectively release their energy in the dense chromosphere. These energetic particles are believed to be responsible (although not always directly) for HXR, EUV, continuum \"white-light\" (WL), infrared and radio emissions. On the basis of temporal correlation, the interconnection between these emissions was early realized \\citep{Na70,Sv70}. Due to the lack of height-resolving observations spectroscopic data have been mostly used to identify the structure of the emitting region \\citep[e.g.][]{Ca74}. Therefore, the actual mechanisms of WL emission and the relation to HXR emitting electrons is still subject of debate \\citep{Ne89,Ne93,2000SoPh..194..305S,Mat03,Po10,Wa10,2011A&A...530A..84K}. It is not even clear whether the emission is optically thick or thin or both. Thus the characteristic heights of continuum WL emission in various models are placed from the photosphere to the upper chromosphere.  Observations supporting both \\citep{1995A&AS..110...99F, Xu06}, a photospheric origin \\citep[e.g.][]{Bo85,1999ApJ...512..454D,Ch06} presumably due to radiative back-warming, and generation of WL emission in the higher temperature regions of the upper chromosphere \\citep[e.g.][]{Ma74, Hu72} have been made.  The major observational challenge here is to have simultaneous HXR, UV and WL images with sufficiently high angular resolution. The relative horizontal positions of HXR and WL flares have been investigated in detail and, although generally been found to agree \\citep[e.g.][]{Kr11, Fl07,Hu06,Me03}, the uncertainty in the relative position was large due to limited absolute pointing accuracy of partial solar disk observations such as {\\it TRACE}. Thus it was not possible to quantify the relative height of the sources above the photosphere in previous studies. With RHESSI \\citep{Li02} data, the height structure of HXR sources has become accessible for detailed observational studies \\citep{As02, Ko08,Ko10, Ba11a, Sa10}. These observations demonstrated that the higher energy HXR footpoints originate from the chromosphere from progressively lower heights, in good agreement with collisional downward propagation of energetic electrons \\citep[see][for details]{Br02}. The newly available SDO continuum images from HMI \\citep{Wa11} and AIA EUV \\citep{Le11} observations allow to scrutinize the spatial structure and more importantly the height structure of various emissions associated with HXR footpoints. In this paper we present the first spatially resolved simultaneous observations of HXR, EUV and WL emission from a footpoint of a solar limb flare observed with RHESSI, SDO/AIA  and HMI. Using the recently developed hard X-ray visibility analysis technique \\citep{Hur02,Sc07} we determine the characteristic heights for the different types of emission using the method presented by \\citet{Ba11a, Koet10, Ko08}. The results are consistent with the EUV and continuum WL being produced above the HXR footpoints by lower energy electrons ($\\sim 12$ keV)  and hence suggest not a photospheric but upper chromospheric origin of WL emission. ",
        "conclusions": "Our observations suggest that the WL emission originates from the upper chromosphere or lower transition region with a density of between  $10^{11}\\mathrm{cm^{-3}}$ and $10^{13}\\mathrm{cm^{-3}}$. These upper and lower limits of the density are due to the uncertainties in the fit of the density scale height. Figure \\ref{aia_height} (bottom) displays the fitted density model, as well as the upper and lower limits for the density. The arrows in Fig. \\ref{aia_height} (top) illustrate the uncertainty in the RHESSI source heights resulting from the uncertainty in the density scale height. The peak of WL emission is found about 1 Mm higher than the position of the lowest observed RHESSI energy range (30 - 40 keV). Note that despite all the uncertainties in the height due to the density fit or due to relative pointing error, the WL source is still above the 30-40 keV RHESSI sources. Due to the presence of the SXR coronal emission in the X-ray spectrum, the non-thermal HXR spectrum below $\\approx$ 20 keV is not known. However, extrapolating the height function found in Section \\ref{srheight} to lower energies, a source at 12 keV would be seen at about 2 Mm. This coincides with the height of the maximum energy deposition as found from Eq. \\ref{eqdudx} for cut-off energies of 12 keV (Fig. \\ref{aiadiff}). The observed position of the maximum of the WL emission would then hint towards energy deposition by the low energy part of the non-thermal electron distribution, contrary to the often made assumption of electron cut-off energies in the range of 20 keV. This also suggest that WL and HXR sources above 30 keV should be spatially separated by $<$ 1 Mm with decreasing separation in events towards the disk center. However, we note that the height of the maximum energy deposition, especially in the case of low initial cut-off energies, is influenced by the loop density and the length of the loop. Assuming that the acceleration of the particles happens in the center of what is observed as the coronal source Fig. \\ref{aiadiff} indicates a precipitation distance of about 13 arcsec from the site of acceleration to the footpoints. This value was used for the computation of the energy deposition rate. In this case, the maximum energy deposition of electrons with energy $\\lesssim 12$ keV will be near the top of the loop. Should the loop density be lower or the acceleration region extended, the lower energy electrons would be depositing their energy deeper. The other interesting aspect of this flare is that HXR emission above $20$~keV is well fitted by a single power-law. Non-uniform ionization of plasma produces a break at the energy corresponding to the column depth of the transition region, which is observed in other flares \\citep{2002SoPh..210..419K,2011ApJ...731..106S}. The absence of such break points to the transition region above the stopping depth of 20 keV electrons again is consistent with the height measurements. The observation of strong emission from heights  $\\ge 2$ Mm suggests that a substantial part of the WL continuum is formed in an optically thin or finite optically thick region with ionized or partially ionized plasma in the upper chromosphere favoring free-bound and free-free emission of the $6173 \\AA$ continuum. This is additionally supported by the height of the AIA emission at 94, 335 and 193 $\\AA$ (representing temperatures $>10^6$ K) originating from around 3 Mm and above the WL emission peak. Finally, it is likely that the WL is powered by lower energy electrons $\\sim 12$~keV, which do not penetrate deep into the solar chromosphere.  \\mdseries"
    },
    "1107/1107.4569_arXiv.txt": {
        "abstract": "\\indent Turbulence plays a major role in the formation and evolution of molecular clouds. The problem is that turbulent velocities are convolved with the density of an observed region. To correct for this convolution, we investigate the relation between the turbulence spectrum of model clouds, and the statistics of their synthetic observations obtained from Principal Component Analysis (PCA). We apply PCA to spectral maps generated from simulated density and velocity fields, obtained from hydrodynamic simulations of supersonic turbulence, and from fractional Brownian motion fields with varying velocity, density spectra, and density dispersion. We examine the dependence of the slope of the PCA pseudo structure function, $\\alpha_\\mathrm{PCA}$, on intermittency, on the turbulence velocity ($\\beta_v$) and density ($\\beta_n$) spectral indexes, and on density dispersion. We find that PCA is insensitive to $\\beta_n$ and to the log-density dispersion $\\sigma_s$, provided  $\\sigma_s$ $\\leq$ 2. For $\\sigma_s$ $>$ 2, $\\alpha_{PCA}$ increases with $\\sigma_s$ due to the intermittent sampling of the velocity field by the density field. The PCA calibration also depends on intermittency. We derive a PCA calibration based on fBms with $\\sigma_s$ $\\leq$ 2 and apply it to 367 \\CO spectral maps of molecular clouds in the Galactic Ring Survey. The average slope of the PCA structure function, $\\langle\\alpha_\\mathrm{PCA}\\rangle=0.62\\pm0.2$, is consistent with the hydrodynamic simulations and leads to a turbulence velocity exponent of $\\langle\\beta_v\\rangle=2.06\\pm0.6$ for a non-intermittent, low density dispersion flow. Accounting for intermittency and density dispersion, the coincidence between the PCA slope of the GRS clouds and the hydrodynamic simulations suggests $\\beta_v$ $\\simeq$ 1.9, consistent with both Burgers and compressible intermittent turbulence. ",
        "introduction": "\\indent Turbulence plays a major role throughout the entire lifetime of a molecular cloud, from its formation to its fragmentation and collapse into clumps, cores, and stars. Supersonic turbulence is not only responsible for most of the kinetic energy supporting clouds against gravity \\citep{williams00},  but is also intrinsically linked to the formation and structure of molecular clouds \\citep[e.g.,][]{maclow04, elmegreen04, audit05, heitsch05, vazquez06, vazquez07, dobbs06, tasker09, banerjee09, glover10, klessen10, federrath11}. Turbulent shocks create local, over-dense regions within molecular clouds that may collapse into sheets, filaments and pre-stellar cores under the effect of self-gravity, a phenomenon known as turbulent fragmentation. In this context, the length scale of the energy source that drives turbulence can explain differences in star-formation efficiency and type --- clustered versus isolated \\citep{klessen00}. In addition, the Initial Mass Function (IMF), which describes the relative probability of stars with different masses when they form, may directly result from the magnitude and spatial scale of turbulent fluctuations \\citep{padoan02, hennebelle08, hennebelle09}. Hence, most analytic models of star formation rely on the index of the turbulent energy spectrum (i.e., the Fourier spectrum of the velocity fluctuations), and on the Probability Distribution Function (PDF) of the turbulent density field \\citep[e.g.,][]{elmegreen97, padoan02, krumholz05, hennebelle08, hennebelle09}.  The energy spectrum is defined as $E({\\bf k})=4\\pi k^2\\left |{\\bf \\hat v}({\\bf k})\\right |^2$, where ${\\bf \\hat v}({\\bf k})$ is the Fourier transform of the velocity field, and obeys a power law ($E(k)$ $\\propto$ $k^{-\\beta_v}$), where $\\beta_v=5/3$ for incompressible (Kolmogorov) turbulence \\citep{frisch95}, and $\\beta_v=2$ for pressureless (i.e., highly supersonic) shock-dominated turbulence, also called Burgers turbulence \\citep{burgers74, passot88}. Intermittency, characterized by fluctuations occurring sporadically (both spatially and temporally) in the turbulent flow, also affects the nature of the turbulent cascade from large to small scales. Numerical simulations show that intermittency is mainly caused by the interaction of strong shocks, causing rare, strong density enhancements. Manifestations of intermittency are observable in the PDF of the density field \\citep{klessen00b, kritsuk2007,schmidt09,fed09b}, and in the index $\\beta_v$ of the energy spectrum \\citep{dubrulle94, she94, boldyrev02a, boldyrev02b, schmidt08}. \\\\ \\indent While Kolmogorov (incompressible) turbulence does not seem appropriate for molecular clouds, in which turbulent motions are highly supersonic (hence compressible), the application of Burgers turbulence is limited to compressive fields for which $\\nabla \\times {\\bf v} = 0$  (e.g., expanding shells). Three-dimensional numerical simulations of both decaying and driven turbulence, however, show that the ratio of compressive to solenoidal (for which $\\nabla \\cdot {\\bf v} = 0$) velocity dispersion is $\\gamma = \\langle{\\bf v_c^2}\\rangle/\\langle{\\bf v_s^2}\\rangle = 0.1$--1 \\citep{porter98, porter99, fed09b}, where ${\\bf v_c}$  and ${\\bf v_s}$ are the compressive and solenoidal parts of the velocity field, respectively. An alternative, which includes intermittency and compressibility effects, was proposed on the basis of the \\citet{she94} model by \\citet{dubrulle94, boldyrev02a, boldyrev02b, schmidt08}. In this theoretical context, the turbulent energy cascade in the inertial range exhibits properties of Kolmogorov turbulence, while close to the dissipative range, intermittent shock structures resulting from compressibility effects start to dominate energy transfer mechanisms. Depending on the dimension of the dominant dissipative, intermittent structures, log-Poisson models predict $\\beta_v = 1.74$--1.83.\\\\ \\indent The type of turbulence (Kolmogorov, Burgers, log-Poisson), along with the density spectrum and density PDF, influence the resulting stellar mass spectrum,  the IMF \\citep[e.g.,][]{padoan02, hennebelle08, hennebelle09}. While the density PDF is better constrained with extinction data \\citep{cambresy99, ossenkopf01, brunt10b}, the density spectrum has been investigated both with FIR dust emission observations \\citep[e.g.,][]{block10} and with \\his 21 cm spectral line observations \\citep{stanimirovic99, elmegreen01}. The index of the energy spectrum, $\\beta_v$, is the subject of many spectral line studies of molecular clouds \\citep{heyer97, brunt02b, ossenkopf02, brunt03, heyer04, hilyblant08}, including this work.  \\citet{heyer97} proposed to use Principal Component Analysis (PCA) to determine the index of the energy spectrum of molecular clouds using molecular line observations. PCA provides pairs of spatial and velocity scales, or ``PCA pseudo structure functions'', which describe the amount of velocity fluctuations contained within an eddy of a given spatial scale. PCA pseudo structure functions thus provide a description of the turbulent energy cascade. Since knowing the turbulence energy spectrum of molecular clouds has great implications for our understanding of fragmentation and star formation, our ultimate goal is to derive the exponent of the velocity power spectrum from the analysis (e.g., via PCA) of spectral line maps. PCA pseudo structure functions, however, are based on emission, and hence on a complex convolution between the density, velocity, and temperature fields, while a complete description of a turbulent gas flow requires separate statistical descriptions of the velocity and density fields. For instance, turbulent velocity fluctuations can make two physically separated emitting elements at different distances from the observer seem to overlap in velocity space \\citep{ballesteros02}. This would produce an element of double the emissivity, which could be interpreted as a single emitting element of double density or column density. Temperature fluctuations also produce emissivity variations that can be interpreted as column density variations. This is not only a limitation of PCA, but also of all observational studies aimed at deriving the energy spectrum from various mathematical tools, such as Centroid Velocity Increments \\citep[see, e.g., ][]{hilyblant08, brunt04}, Velocity Channel Analysis \\citep{lazarian00, esquivel03}, and Velocity Coordinate Spectrum \\citep{lazarian06}. Unfortunately, while observations provide column densities and their PDF, the temperature, velocity, and volumetric density fields, however, are not directly observable. In order to determine the relation between the PCA pseudo structure function and the Fourier spectrum of the velocity field, we therefore need to disentangle the density, velocity, and temperature contributions to the observed position-position-velocity data and to the shape of the PCA pseudo structure function. This can be done by comparing the PCA pseudo structure function of simulated spectral maps obtained from simulations of molecular clouds to the statistics of the input velocity, density, and temperature fields. Here, we use isothermal simulations, and hence neglect the molecular excitation problem and the effects of star formation on the gas temperature to concentrate on density and velocity fluctuations.  Thus, our goal is to establish a {\\it calibration} relation between the slope of the PCA pseudo structure function and the slope of the turbulence velocity spectrum, and to investigate how this calibration relation varies with properties of the density field. \\\\ \\indent \\citet{brunt03etal} established a calibration relation based on MHD simulations. Their simulations included many of the important physical processes in the ISM, such as gravity, magnetic fields, star formation feedback, heating and cooling.  The focus of their study was to link the PCA derived relationship between velocity differences and spatial scale to a particular order of structure function. They were able to show that PCA pseudo structure functions correspond to a low order ($\\simeq$ 1) structure function even in the regime of strong intermittency. The resolution of their simulation was however too small to allow the existence of an inertial range, where the slope of the power-law velocity spectrum can be measured. As a result, their simulation had to be modified a posteriori in Fourier space to create a power-law velocity spectrum. In addition to this limitation, the turbulence in the MHD simulations by \\citet{brunt03etal} were mostly driven compressively by expanding shells, occurring in star-forming regions.\\\\ \\indent  In this paper, we investigate the statistical relation between the PCA pseudo structure function and the intrinsic index of the velocity field for two types of forcing --- compressive and solenoidal --- based on hydrodynamic simulations with a distinct inertial range \\citep{fed09b}, combined with Fractional Brownian Motion (fBm) simulations. We examine the sensitivity of this relation to the statistics of the density field by varying both the exponent of the power-law density spectrum and the density PDF of the fBm simulations. Section~\\ref{pca_section} describes PCA and the details of the method. Section~\\ref{simulations_section} describes the simulations. In Sections~\\ref{pca_hydro_section} and \\ref{pca_fbm_section}, we present the results of PCA applied to the hydrodynamic and fBm simulations, and provide a PCA calibration in Section~\\ref{calibration_lognormal_section}. We apply PCA to spectral observations of molecular clouds taken as part of the Galactic Ring Survey \\citep{GRS} in Section~\\ref{application_section}. Section~\\ref{conclusion} consist of a brief conclusion. \\\\ ",
        "conclusions": "\\indent We applied Principal Component Analysis (PCA) to synthetic Position-Position-Velocity (PPV) spectral maps generated from the density and velocity fields of solenoidally and compressively forced hydrodynamic simulations of supersonic turbulence \\citep{fed08, fed09a, fed09b}, and of fractional Brownian motion simulations, in order to constrain the calibration relation between the PCA pseudo structure function and the index of the velocity spectrum of turbulence, and to examine the dependency of this relation on the density spectrum, intermittency, and density dispersion.\\\\ \\indent We demonstrated that the calibration relation, the relation between the slope of the PCA structure function $\\alpha_{PCA}$ and $\\beta_v$, does not depend on the exponent of the power-law density spectrum $\\beta_n$. \\\\ \\indent For a log-density dispersion $\\sigma_s$ $\\leq 2$,  we do not find any dependence of the PCA calibration on the dispersion of the density PDF. We derive a PCA calibration relation, $\\beta_v = 0.20\\pm0.05 + (2.99\\pm0.09)\\alpha_{PCA}$ valid for $\\sigma_s$ $\\leq$ 2 and $\\beta_v$ $=$ 1.2--2.6. For $\\sigma_s$ $>$ 2, we find a strong dependence of the calibration between $\\alpha_{PCA}$ and $\\beta_v$ with $\\sigma_s$. Extreme density fluctuations intermittently sample the velocity field, producing an effect similar to intermittency in the velocity field itself - i.e. mimicking discontinuous velocity jumps, although the detailed mechanism is rather different. PCA is stable below a threshold of the log-density dispersion, $\\sigma_s$ $\\simeq$ 2, but if real molecular clouds exceed this, then an additional overestimation factor applies to $\\alpha_{PCA}$. Without knowledge of the true 3D log-density dispersions in the cloud sample, the estimation of the turbulent spectrum in molecular clouds remains uncertain. \\citet{brunt10a, brunt10b} developed a method to estimate the density PDF of molecular clouds based on the 2D power spectrum, the variance, and the PDF of the 2D column density, from which the 3D density PDF can be reconstructed, even in cases where the density PDF is not lognormal. However, this method requires high fidelity measures of column density such as extinction derived from 2MASS photometry of background stars and high spatial dynamic range. Therefore, it is not readily applicable to our set of data from the Galactic Ring Survey for which the spatial dynamic range for most clouds is limited. In addition, numerical simulations predict a relation between the log-density dispersion and the Mach number \\citep{price11}, but this relation also depends on the relative contribution of solenoidal and compressive modes \\citep{fed09b}. An initial test of the log- density dispersion - Mach number relation has been made \\citep{brunt10a, brunt10b}, and this suggests that both solenoidal and compressive forcing are important and that density dispersions are likely to be high enough that their effect on PCA is not insignificant. \\\\ \\indent We demonstrated that intermittency in the velocity field also increases the PCA slope for a given velocity spectrum. This effect is due to a combination of the operating order of PCA (PCA traces the first order structure function exponent, $\\zeta(1)$), and the variation of the ratio between $\\zeta(1)$ and $\\beta_v$ with intermittency. Thus, if a first-order scheme is used to measure the second order exponent $\\beta_v$, then some knowledge of the level of intermittency is required. By accounting for the level of intermittency, we were able to reconcile PCA measurements between non-intermittent fBms and the HD fields.  \\\\ \\indent We applied PCA to \\CO spectral maps of 367 molecular clouds identified in the Galactic Ring Survey \\citep{GRS}. We found that the average slope of the PCA pseudo structure function and the slope of the composite structure function, made of all the spatial and velocity scales derived in all the GRS clouds, are consistent with $\\alpha_{PCA}$ $=$ $0.62\\pm0.2$. Applying the PCA calibration obtained from fBms with $\\sigma_s$ $\\leq$ 2, the PCA slope obtained for GRS molecular clouds corresponds to an average turbulence spectral index of $\\langle\\beta\\rangle=2.06\\pm0.6$. However, we have shown that intermittency and density dispersion need to be taken into account.  The average PCA slope obtained for the GRS clouds is in very good agreement with the PCA slope obtained from both solenoidally and compressively forced HD simulations, albeit in better agreement  (at $<$ 1 $\\sigma$) with the solenoidally forced HD simulations. This agreement suggests that  turbulence in molecular clouds, as in the HD simulations, obey log-Poisson scaling relations (intermittent, compressible turbulence) with 2D shocks as the dominant dissipative structures."
    },
    "1107/1107.6046_arXiv.txt": {
        "abstract": "Using Hubble Space Telescope (HST) photometry, we characterize the age of the stellar association in the vicinity of supernova (SN) 2011dh and use it to infer the zero-age main sequence mass (\\mzams) of the progenitor star.  We find two distinct and significant star formation events with ages of $<$6 and 17$^{+3}_{-4}$~Myrs, and the corresponding \\mzams\\ are $>$29 and 13$^{+2}_{-1}$~\\msun, respectively.  These two bursts represent 18$^{+4}_{-9}$\\% (young) and 64$^{+10}_{-14}$\\% (old) of the total star formation in the last 50~Myrs.  Adopting these fractions as probabilities suggests that the most probable \\mzams\\ is $13^{+2}_{-1}$~\\msun. These results are most sensitive to the luminosity function along the well-understood main sequence and are less sensitive to uncertain late-stage stellar evolution. Therefore, they stand even if the progenitor suffered disruptive post-main-sequence evolution (e.g. eruptive mass loss or binary Roche-lobe overflow).  Progenitor identification will help to further constrain the appropriate population.  Even though pre-explosion images show a yellow supergiant (YSG) at the site of the SN, panchromatic SN light curves suggest a more compact star as the progenitor.  In spite of this, our results suggest an association between the YSG and the SN. Not only was the star located at the SN site, but reinforcing an association, the star's bolometric luminosity is consistent with the final evolutionary stage of the 17~Myr old star burst. If the YSG disappears, then \\mzams$=13^{+2}_{-1}$~\\msun, but if it persists, then our results allow the possibility that the progenitor was an unseen star of $>$29~\\msun. ",
        "introduction": "\\label{section:introduction}  Observational measurements of the masses of supernova (SN) progenitors are currently scant and highly uncertain.  Of the $\\sim$20 SNe that have progenitor mass constraints, only about half have well-defined masses and the rest have upper bounds \\citep{arnett89,aldering94,barth96, vandyk99,vandyk02,smartt02,vandyk03a,vandyk03b, smartt04,maund05, hendry06,li05,li06,li07,galyam07,crockett08,galyam09,smartt09a, smartt09b,smith11b,smith11c}. In general, stellar evolution theory predicts a clear mapping between the zero-age main sequence mass (\\mzams) and the explosion scenario for isolated massive stars \\citep{woosley02,heger03a,dessart11}. These limited observations suggest that the least massive stars explode as SN II-P as expected \\citep{smartt09a}. More massive stars are expected to lose much of their mass and explode as H-deficient SNe (IIb and Ib/c). However, some H-rich SNe (in particular IIn) have been associated with very massive stars \\citep{galyam09,smith11b,smith11c}. Furthermore, theory \\citep{claeys11,dessart11} and the relatively high observed rates of H-deficient SNe \\citep{smith11c} imply that binary evolution may figure prominently in producing the H-deficient SNe. Even in light of these complications, the maximum \\mzams\\ associated with SN II-P is much lower than expected \\citep{smartt09a}.  While tantalizing, these initial results are poorly constrained, and even the simple statement that stars more massive than $\\sim$8~\\msun\\ explode as SNe requires more observational constraints.  Thus every new supernova marks an important opportunity to add a new progenitor mass measurement to this poorly-constrained sample.  Recently, amateurs and the Palomar Transient Factory (PTF) collaboration detected supernova (SN) 2011dh in M51, the Whirlpool Galaxy \\citep{griga11,silverman11,arcavi11b}. Initial reports classified it as II \\citep{silverman11}, but late time spectra reveals this explosion to belong to the rare IIb transitional class \\citep{arcavi11b,marion11}, indicating that the progenitor lost much of its hydrogen envelope due to either binary evolution or unknown single-star mass-loss physics. The close proximity of this SN \\citep[$(M-m)_0=29.42$, 7.7 Mpc;][companion galaxy NGC~5195]{tonry01} and the wealth of Hubble Space Telescope (HST) data on M51 offers a rare opportunity to further explore the physics of the progenitor.   Using archival HST photometry, \\citet{maund11} and \\citet{vandyk11} identified a yellow supergiant (YSG) as the progenitor candidate.  \\citet{vandyk11} fit the magnitude and color of the progenitor to evolutionary tracks and derive a zero-age-main-sequence progenitor mass of \\mzams = 18-21~\\msun. However, \\citet{maund11} argue that an uncertain mass-loss history causes the color to be an unreliable characteristic of the progenitor. Instead, they treat the bolometric luminosity as an intrinsic property determined by the core mass, which is in turn determined by \\mzams.  Matching the bolometric luminosity to the very last stages of evolutionary models, they derive \\mzams = 13$\\pm 3$~\\msun.  To further complicate the situation, recent observations suggest that the YSG may not even be the progenitor.  Based upon the characteristics of the SN light curves, \\citet{arcavi11b} and \\citet{soderberg11} argue that the YSG is not the progenitor and instead suggest a more compact source.  This discrepancy in mass measurements highlights the fact that interpretation of the precursor photometry alone is sensitive to uncertain mass-loss physics in the final evolutionary stages of the models.  Furthermore, using precursor imaging to measure progenitor masses requires that precursor imaging exists and that the SN position is known to sub-arcsecond precision.  The majority of past SNe have neither pre-existing HST imaging, nor sufficiently accurate astrometry.  A significant fraction of future nearby SNe will lack precursor imaging as well.  In this letter, we address the progenitor mass with an independent approach; we characterize the age of the stellar association surrounding SN~2011dh and derive the corresponding \\mzams\\ for the progenitor.  This technique was thoroughly described by \\citet{gogarten09} on a SN ``impostor'' in NGC~300 and by \\citet{badenes09} on SN remnants in the Magellanic Clouds.  In brief, using well-established stellar population modeling techniques, we can age-date the star formation episode that led to the observed SN.  The resulting age can place strong constraints on the mass of the precursor by leveraging the well-understood properties of a large number of main-sequence stars. This complementary approach is feasible even when there is no precursor imaging, or when the SN position is only localized to within several arcseconds.  In \\S2, we detail our application of this technique to SN~2011dh, the most distant and only ongoing SN to which it has been applied.  In \\S3, we provide the resulting age and mass measurements, and in \\S4, we compare our resulting progenitor mass with those obtained using direct imaging of the progenitor \\citep{maund11,vandyk11}.  We find our results to be most consistent with the results of \\citep{maund11} and improve upon their uncertainties. ",
        "conclusions": "\\label{section:conclusions}  Using HST photometry, we characterize the age of the stars in the vicinity of SN 2011dh, and from this age, we infer the \\mzams\\ of the progenitor star.  The recent SFH shows two SF events with ages of $<$6 and 17$^{+3}_{-4}$~Myrs, corresponding to \\mzams\\ for the progenitor of $>$29 and 13$^{+2}_{-1}$~\\msun, respectively.  Given that the older burst at 17$^{+3}_{-4}$~Myr represents 64$^{+10}_{-14}$\\% of the total SF in the last 50~Myrs, the most probable \\mzams\\ for the progenitor is 13$^{+2}_{-1}$~\\msun.  Because MATCH leverages the entire CMD, including the MS, these \\mzams estimates are relatively insensitive to the individual peculiarities of post-MS or binary evolution.  Stellar population analysis shows that the colors of the YSG are consistent with the surrounding stars only if it experienced peculiar evolution.  Testing the robustness of the three SF events, we perform several jackknifing tests, removing a bright star and finding new best-fit SFHs.  The $<$6 and 17~Myr old populations are robust to these tests, but upon removing the YSG, the 8~Myr old burst vanishes ($<$3.5\\% of the recent SF).  Hence, the 8~Myr feature is solely dependent upon the color and magnitude of the most poorly modeled star in the field, the YSG -- an unlikely scenario.  More likely, the YSG was born in one of the other events and experienced disruptive mass loss, giving it peculiar colors. In fact, its bolometric luminosity is consistent with the final evolutionary stage of the 17~Myr old population.   Regardless of an association with the SN, the luminosity of the YSG \\citep{maund11,vandyk11} suggests that it is associated with the 17~Myr population, in which case, we derive a \\mzams\\ for the YSG of 13$^{+2}_{-1}$~\\msun.  The previous \\mzams\\ estimates were obtained by comparing photometry of the YSG with stellar evolution models.  Of the two previous attempts to estimate \\mzams\\, one at 13$\\pm$3~\\msun\\ \\citep{maund11} and one at 18-21~\\msun\\ \\citep{vandyk11}, the lower mass estimate is most consistent with our results.  This consistency leads to a couple of conclusions.  For one, this consistency validates using nearby stellar ages to estimate the progenitor mass when an archival HST image of the progenitor is not available.  Given the distance to M51, $\\sim$7.7~Mpc, the fact that this technique is able to further constrain the \\mzams\\ estimate is remarkable.  Directly modeling the progenitor in precursor imaging only requires accurate photometry for the most luminous stars and has become a routine exercise up to 20~Mpc \\citep{smartt09b}.  Hence, for direct imaging of the precursor, SN~2011dh is quite close.  Our technique, on the other hand, relies on having a sufficient number of main sequence stars to characterize the age of the associated star burst. With a distance of $\\sim$7.7~Mpc, the magnitude limit is quite high, resulting in only 16 detectable upper main sequence stars within 50 pc of the SN.  Even with this small number of upper main sequence stars, we were able to improve upon the direct imaging mass constraint for the YSG.  The second conclusion we draw from the consistency is that fitting the bolometric luminosity of the progenitor to the last stages of evolutionary tracks \\citep{maund11} is a more robust method to estimate \\mzams\\ than fitting both the luminosity and effective temperature \\citep{vandyk11}.  The literature is full of theoretical arguments suggesting that the bolometric luminosity during the last evolutionary stages is a more intrinsic property of \\mzams\\ \\citep{arnett89,woosley02,smartt09b}, but as far as we know, there have been no independent observations to support this.  Because age-dating the nearby stars leverages information from the MS phase, our results provide a complimentary estimate of \\mzams\\ and independent support for using the bolometric luminosity only in direct imaging techniques.  While the most probable \\mzams\\ of the SN progenitor is 13$^{+2}_{-1}$~\\msun, there is a small but not insignificant probability that the progenitor mass is $>$29~\\msun.  Ruling out one or the other as the progenitor requires further information. If the YSG persists in post-SN imaging, then the progenitor was indeed compact \\citep{arcavi11b,soderberg11}, and neither mass estimate is ruled out. However, if the YSG disappears, then our progenitor mass estimate of 13$^{+2}_{-1}$~\\msun\\ would validate and further constrain the \\citet{maund11} estimate. Given the YSG's positional coincidence with the SN and the fact that its bolometric luminosity is consistent with the final evolutionary stage of the 17~Myr old population, it would be odd if the YSG is not associated with the SN."
    },
    "1107/1107.4872_arXiv.txt": {
        "abstract": "Charged-current neutrino-nucleus cross sections for $^{54,56}$Fe and $^{58,60}$Ni are calculated and compared using frameworks based on relativistic and Skyrme energy density functionals, and the shell model. The current theoretical uncertainties in modeling neutrino-nucleus cross sections are assessed in relation to the predicted Gamow-Teller transition strength and available data, multipole decomposition of the cross sections, and cross sections averaged over the Michel flux and Fermi-Dirac distribution. Employing different microscopic approaches and models, the DAR neutrino-$^{56}$Fe cross section and its theoretical uncertainty are estimated: $<\\sigma>_{th}$=(258$\\pm$57) $\\times$10$^{-42}$cm$^2$, in very good agreement with the experimental value: $<\\sigma>_{exp}$=(256$\\pm$108$\\pm$43) $\\times$ 10$^{-42}$cm$^2$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.4275_arXiv.txt": {
        "abstract": "We analyse how the structure of the inner curst is influenced by the pairing correlations. The inner-crust matter, formed by nuclear clusters immersed in a superfluid neutron gas and ultra-relativistic electrons, is treated in the Wigner-Seitz approximation. The properties of the Wigner-Seitz cells, i.e., their neutron to proton ratio and their radius at a given baryonic density, are obtained from the energy minimization at beta equilibrium. To obtain the binding energy of baryonic matter we perform Skyrme-HFB calculations with zero-range density-dependent pairing forces of various intensities. We find that the  Wigner-Seitz cells have much smaller numbers of protons compared to previous calculations. For the dense cells the binding energy of the configurations with small proton numbers do not converge to a well-defined minimum value which precludes the determination of their structure. We show that for these cells there is a significant underestimation of the binding energy due to the boundary conditions at the border of the cells imposed through the Wigner-Seitz approximation. ",
        "introduction": "The inner-crust of neutron stars extends from the so-called neutron drip density, $\\rho_d \\approx 4 \\times 10^{11}$ g cm$^{-3}$, defined as the density where the neutrons start to drip out from the nuclei of the crust, up to a density of about $\\rho \\approx 1.4 \\times 10^{14}$ g cm$^{-3}$ at which there is a transition towards the uniform core matter. The size and the precise density limits of the inner crust depend on the star mass and on the equation of state employed in the star models \\cite{glendenning,haensel_book}.   The inner-crust matter of non-accreting cold neutron stars is most probably formed by a crystal lattice of nuclear clusters immersed in a sea of low-density superfluid neutrons and ultra-relativisitic electrons. It is generally considered that in most of its part the inner-crust is formed of nuclei-like clusters. More complex \"pasta\" structures (e.g., rods, plates, bubbles) are expected to be formed in the transition region between the inner crust and the core matter, see for instance Refs.~\\cite{glendenning,haensel_book} and references therein.   The first microscopic calculation of the inner-crust structure, still used as a benchmark in neutron stars studies (e.g., Refs.~\\cite{page,pizzochero,monrozeau}) was performed by Negele and Vautherin in 1973~\\cite{negele}. In this work the crystal lattice is divided in spherical cells which are treated in the Wigner-Seitz (WS) approximation. The nuclear matter from each cell is described in the framework of Hartree-Fock (HF) approximation based on the Density Matrix Expansion (DME)~\\cite{negele72}. This approach was preferred to the Density-Dependent Hartree-Fock theory~\\cite{negele70} in order to reduce computational complication induced by the non-local exchange potential. The parameters of the DME theory were adjusted to reproduce the experimental binding energies of atomic nuclei and the theoretical calculation of infinite neutron matter available at that time. The spin-orbit interaction was taken into account for the protons but neglected for the neutrons. The HF equations were solved in coordinate representation imposing mixed Dirichlet-Neuman boundary conditions at the border of the cells. The properties of the WS cells found in Ref.~\\cite{negele}, determined for a limited set of densities, are shown in Table~\\ref{tab:NV}. The most remarkable result of this calculation is that the majority of the cells have semi-magic and magic proton numbers, i.e., Z=40,50. This indicates that in these calculations there are strong proton shell effects, as in isolated atomic nuclei. However, as seen from Fig.~\\ref{fig:gap} of Ref.~\\cite{negele} the  energies corresponding to the cells configurations based on various Z numbers are in fact very close to each other. This fact raises several questions: i) what is the sensitivity of the results on the nuclear interaction, ii) how much the  pairing correlations influence the cells structure, iii) how reliable is the WS approximation.   The effect of pairing correlations on the structure of WS cells was investigated in Ref.~\\cite{baldo} within the HF+BCS approach. In the most recent version of these calculations the authors solved the HF+BCS equations with a mixture between the phenomenological functional of Fayans et al~\\cite{fayans}, employed in the nuclear cluster region, and a microscopical functional derived from Bruckner-Hartree-Fock calculations in infinite neutron matter. The latter was used to describe the neutron gas in the WS cells. In this framework it was found that the cells have not a magic or semi-magic number of protons, as in Ref.~\\cite{negele}. It was also found that pairing can change significantly the structure of the cells compared to HF calculations. These findings show that in order to determine the most probable structure of the inner crust one needs more investigations based on various effective interactions and many-body approximations.  In this study we analyse the effect of pairing on inner crust structure in the Hartree-Fock-Bogoliubov approach (HFB). This approach offers better grounds than HF+BCS approximation for treating pairing correlations in non-uniform nuclear matter with both bound and unbound neutrons. To investigate the dependence of the inner crust structure on pairing, the HFB calculations are performed with three different density-dependent pairing forces adjusted to reproduce various pairing scenarios in nuclear matter.  In principle, the symmetries of the inner crust lattice should be taken properly into account when inner crust structure is determined. Since imposing the exact lattice symmetries in microscopic self-consistent models is a very difficult task (for approximative solutions to this problem see ~\\cite{chamel2010} and the references therein) we solve the HFB equations in the WS approximation, as commonly done in inner crust studies \\cite{negele,baldo}. This approximation induces, through the boundary conditions at the border of the cells, an artificial shell structure in the  energy spectra of nonlocalized  neutrons \\cite{baldo2006}. The errors caused by the spurious shells, which affect mainly the high density cells, are estimated  by using the method proposed in Ref.\\cite{margueron2007}.  The paper is organized as follows: in Section II we describe how the cells energy is calculated in the HFB and WS approximations, in Section III we present the equations for beta equilibrium and in Section IV we discuss the results of the calculations. ",
        "conclusions": ""
    },
    "1107/1107.1666_arXiv.txt": {
        "abstract": "The Transit Timing Variations (TTVs) technique provides a powerful tool to detect additional planets in transiting exoplanetary systems. In this paper we show how transiting planets with significant TTVs can be systematically missed, or cataloged as false positives, by current transit search algorithms, unless they are in multi-transit systems. If the period of the TTVs, $P_{TTV}$, is longer than the time baseline of the observations and its amplitude, $A_{TTV}$, is larger than the timing precision limit of the data, transiting planet candidates are still detected, but with incorrect ephemerides. Therefore, they will be discarded during follow-up. When $P_{TTV}$ is shorter than the time baseline of the observations and $A_{TTV}$ is sufficiently large, constant period search algorithms find an average period for the system, which results in altered transit durations and depths in the folded light curves. Those candidates can get subsequently discarded as eclipsing binaries, grazing eclipses, or blends. Also, for large enough $A_{TTVs}$, the transits can get fully occulted by the photometric dispersion of the light curves. These detection biases could explain the observed statistical differences between the frequency of multiple systems among planets detected via other techniques and those detected via transits. We suggest that new transit search algorithms allowing for non-constant period planets should be implemented. ",
        "introduction": "In a transiting exo-planetary system, the presence of other planets can be inferred by the effect of their gravitational pull on the transiting planet. That effect can be detected as periodic transit timing variations (TTVs), with amplitudes potentially as large as several tens of minutes \\citep[e.g][]{Miralda2002, Agol2005, Holman2005, Heyl2007, Veras2010}.  These theoretical predictions have been confirmed by two recent results from the Kepler satellite mission \\citep{Borucki2004}. In the first of those results, the system Kepler- 9, composed of two transiting Saturn-size planets, presents TTVs with average amplitudes of 4 and 39 minutes per orbit \\citep{Holman2010}. The other reported system, Kepler-11, composed of five $2.3 - 13.5 M_{Earth}$ transiting planets within 0.25AU from the star, and a sixth more massive planet in an outer orbit, shows TTVs amplitudes of up to 30 minutes \\citep{Lissauer2011}. \\citet{Steffen2010} have also published five Kepler systems with more than one transiting object, although those have not been confirmed as planets.  It is surprising that TTVs have been first detected in multiple transit planet systems, i.e. systems where several planets are co-planar within a few degrees, since those are geometrically less likely than non-coplanar multiple planet systems where only one planet transits. However, no definitive detection of single-transiting planets TTVs has been reported. Some authors have reported hints of possible TTVs, for WASP-3b \\citep{Maciejewski2010}, WASP-10b \\citep{Maciejewski2011}, WASP-5b \\citep{Fukui2011}, and most recently HAT-P-13b \\citep{Nascimbeni2011,Pal2011}, but none of those claims  has yet been confirmed. \\citet{Ford2011} also recently reported over 60 TTVs candidate systems in the Kepler mission data.  The apparent scarcity of single-transit multi-planet systems was also noticed by \\citet{Fabrycky2009}, who showed how one third of the non-transiting planets discovered via the radial velocity technique, and with orbital periods similar to the transiting planets known at the time, had companions versus no detected companions for any of the transiting planets. Two single-transit planetary systems with a second companion have been detected since then, CoRoT- 7b \\citep{Leger2009, Queloz2009} and HAT-P-13b \\citep{Bakos2009}, but the statistics still disagree.  In this work we explore potential observational and transit search systematics that could introduce single-transit multi-planet systems detection biases. Section 2 summarizes how the transit search algorithms used by successful surveys work and their potential detection biases. Section 3 describes our model simulation for a generic single-transiting system with TTVs and its main observational signatures. Section 4 describes the model simulation results. We discuss our findings in section 5. ",
        "conclusions": "\\label{sec:discussion}  \\citet{Terquem2007} and \\citet{LoCurto2010} have suggested that the absence of multiple planet systems (and therefore TTVs) among transiting planets discovered by surveys (in their majority hot Jupiters discovered from the ground until Kepler was launched) could be due to migration mechanisms. Massive short period planets would have moved inwards disturbing the orbits of all other planets in their way. However, the recent detection with Kepler of over 130 candidate systems showing some evidence of TTVs challenges those hypotheses \\citep{Ford2011}. Figure 3 represents the frequency of TTV candidates reported by \\citet{Ford2011} in two separate histograms; one for planets with $R_p < 6 R_{Earth}$, i.e. Neptune-size and smaller planets, and the other for planets with $6 R_{Earth} < R_p < 15 R_{Earth}$, i.e. gas giants. Given photometric precision and observing cycle limitations, gas giants with orbital periods of six days or less are the planets most likely to be detected from the ground. The bottom diagram in the figure shows five TTV candidates that match those criteria. Assuming all of the \\citet{Ford2011} candidates truly are planets with TTV variations, and comparing those numbers with the total number of hot Jupiters candidates with $P < 6$ days in the current Kepler mission's candidate list ($\\sim 65$ planets), 1 out of each 15 transiting hot Jupiters should  show TTVs. Based on this estimate and given the number of hot Jupiter transiting planets currently confirmed, about 7 systems with significant TTVs should have been already discovered from the ground. Doing a similar estimation with the number of multiple planet systems found to date via the radial velocity technique (source http://exoplanet.eu), we find that 14 have planets with orbital periods $P < 6$ days and $Msini > 0.5 M_J$. Of those 2 are Jupiter-mass planets (assuming $Msini > 0.5 M_J$), which gives a fraction of close-in massive planets in multiple systems of about 1/7, although not all will present significant TTVs . Assuming also that a fraction of the \\citet{Ford2011} TTV candidates will be false detections, we conclude that the true number of hot Jupiters showing TTVs falls somewhere below 1/15, and those systems are most likely being missed by current planet search algorithms.   Our simulations reveal two effects that may explain why transiting planets with significant TTVs elude detection. One is finding the wrong transit ephemeris when $P_{TTV} > P_{obs}$. This effect does not preclude those systems from being labeled as transit candidates, but if a dubious candidate around, for instance, a relatively dim star is detected by a small telescope with $\\sigma_{ph} \\gtrsim 0.01$, and cannot be confirmed by observing it at a higher photometric precision because the ephemerides are incorrect, the candidate will be discarded before spectroscopic confirmation follow-up.  The other effect, i.e. wider, shallower and \"V\"-shaped transits, may explain why planets with $P_{TTV} < P_{obs}$ and significantly large $A_{TTV}$ are lost. Figure 4 shows how, for our simulated 1$R_{jup}$ planet, the depth and length of the transit changes in the folded light curve when $A_{TTV}$ varies from 0.0025 to 0.02. The almost vertical line in the figure indicates, for comparison, how the transit length and depth of a real planet would vary when its radius increased from 1 to 2 $R_{jup}$. Clearly, a modified BLS algorithm , or visual inspection, will find the depth and length of the candidate to be incompatible with the expected parameters of a real planet, and the system will be rejected as a grazing eclipsing binary. In the case of very low light curve scatter and continuous space-borne observations, like Kepler's, TTVs will be more difficult to miss, as show Figures 1 and 2. However, faint targets will encounter the same problems as ground-based observations.  Algorithms like the ones developed by \\citet{Regulo2007} and \\citet{Renner2008} should have been able to detect TTVs in the CoRot data, but the limited monitoring per field \\citep[$\\sim$ 150 days,][]{Auvergne2009}, and strong data systematics may have prevented such detections.  We conclude that, if our hypotheses are correct, the implementation of new transit search algorithms optimized for the detection of TTVs, and their application to existing ground-based transit-search databases, should lead to the detection of new single-transit systems with significant TTVs."
    },
    "1107/1107.2759_arXiv.txt": {
        "abstract": "We present the evolutionary properties and luminosity functions of the radio sources belonging to the Chandra Deep Field South VLA survey, which reaches a flux density limit at 1.4 GHz of $43~\\mu$Jy at the field center and redshift $\\sim 5$, and which includes the first radio-selected complete sample of radio-quiet active galactic nuclei (AGN). We use a new, comprehensive classification scheme based on radio, far- and near-IR, optical, and X-ray data to disentangle star-forming galaxies from AGN and radio-quiet from radio-loud AGN. We confirm our previous result that star-forming galaxies become dominant only below 0.1 mJy. The sub-millijansky radio sky turns out to be a complex mix of star-forming galaxies and radio-quiet AGN evolving at a similar, strong rate; non-evolving low-luminosity radio galaxies; and declining radio powerful ($P \\ga 3 \\times 10^{24}$ W~Hz$^{-1}$) AGN. Our results suggest that radio emission from radio-quiet AGN is closely related to star formation. The detection of compact, high brightness temperature cores in several nearby radio-quiet AGN can be explained by the co-existence of two components, one non-evolving and AGN-related and one evolving and star-formation-related. Radio-quiet AGN are an important class of sub-millijansky sources, accounting for $\\sim 30\\%$ of the sample and $\\sim 60\\%$ of all AGN, and outnumbering radio-loud AGN at $\\la 0.1$ mJy. This implies that future, large area sub-millijansky surveys, given the appropriate ancillary multi-wavelength data, have the potential of being able to assemble vast samples of radio-quiet AGN by-passing the problems of obscuration, which plague the optical and soft X-ray bands. ",
        "introduction": "The relationship between star formation and AGN in the Universe is one of the hottest topics of current extragalactic research, at two different levels. On cosmological scales, the growth of supermassive black holes in AGN appears to be correlated with the growth of stellar mass in galaxies \\citep[e.g.,][]{mer08}. On nuclear scales, the accreting gas feeding the black hole at the center of the AGN might trigger a starburst. The black hole, through winds and jets, can in turn feed energy back to its surroundings, which can compress the gas and therefore accelerate star formation but can also blow it away, thereby stopping accretion and star formation altogether. The general consensus is that nuclear activity plays a major role in the co-evolution of supermassive black holes and galaxies through the so-called ``AGN Feedback'' and indeed radio emission from AGN has been recently suggested to play an important role in galaxy evolution \\citep{cro06}. Moreover, radio observations afford a view of the Universe unaffected by the absorption, which plagues observations made at most other wavelengths, and therefore provide a vital contribution to our understanding of this co-evolution. These two points imply that studies of the evolution of star-forming galaxies (SFG) and AGN in the radio band should provide a better understanding of the link between the two phenomena. These are obviously done best by reaching relatively faint ($\\la 1$ mJy) flux densities, and hence the importance of characterizing the radio faint source population.  After years of intense debate, the contribution to the sub-millijansky population from synchrotron emission resulting from relativistic plasma ejected from supernovae associated with massive star formation in galaxies appears not to be overwhelming, at least down to $\\sim 50~\\mu$Jy, contrary to the (until recently) most accepted paradigm. Our deep ($S_{\\rm 1.4GHz} \\ge 43~\\mu$Jy) radio observations with the NRAO Very Large Array (VLA) of the Chandra Deep Field South (CDFS), complemented by a variety of data at other frequencies, imply in fact a roughly 50/50 split between SFG and AGN \\citep{pad09}, in broad agreement with other recent papers \\citep[e.g.,][]{sey08,smo08}.  The purpose of this paper is to study the evolution and luminosity functions (LFs) of sub-millijansky radio sources through the VLA-CDFS sample. Apart from the topics mentioned above, this is important also for other issues, including:  \\begin{enumerate}  \\item predictions for the source population at radio flux densities $< 1~\\mu$Jy, which are relevant, for example, for the Square Kilometre Array (SKA). All existing estimates, in fact, had to rely, for obvious reasons, on extrapolations and are based on high flux density samples. This affects particularly the highest redshifts, which can better be probed at fainter flux densities;  \\item the radio evolution of radio-quiet AGN. No radio-selected sample of radio-quiet AGN is currently available and this is badly needed to shed light on the mechanism behind their radio emission and allow a proper comparison with radio-loud quasars;  \\item the fact that number counts by themselves do not necessarily reflect the relative intrinsic abundance of astrophysical sources, which requires the determination of the evolution and LF \\citep{pad07}.  \\end{enumerate}  We note that the evolution and LFs of sub-millijansky radio sources have been studied, so far, only in two fields: the COSMOS field \\citep{smo09a,smo09b} and the Deep {\\it Spitzer} Wide-area InfraRed Extragalactic (SWIRE) field \\citep{str10}, in both cases up to a maximum redshift of $1.3$ and without differentiating between radio-quiet and radio-loud AGN. Source classification was based on a rest-frame optical color scheme and on spectral energy distribution (SED) fitting to photometric data covering the UV to near-IR range respectively.  We define as AGN sources in which most of the energy is produced through physical processes other than the nuclear fusion that powers stars. In practice, this means that electromagnetic emission most likely related, directly or indirectly, to a supermassive black hole is predominant in at least one band. A small fraction of AGN have, for the same optical power, radio powers three to four orders of magnitude higher than the rest. These are called ``radio-loud'' quasars and most of the energy they emit is non-thermal and is associated with powerful relativistic jets, although thermal components associated with an accretion disk may also be observed, especially in the optical/UV band. Radio galaxies are also characterized by strong radio jets (manifested also through radio lobes), typically laying in or near the plane of the sky, and a fraction of them (the most powerful ones) are thought to be radio-loud quasars, which have instead their jets oriented with a small angle to the line of sight \\citep[e.g.,][]{up95}. We define as ``radio-quiet'' AGN in which jets are either not present or make a tiny contribution to the total energy budget over the whole electromagnetic spectrum, which is dominated by thermal emission. All other AGN we call radio-loud. Note that radio-quiet AGN are not radio-silent. Indeed, the radio power of many low-luminosity radio galaxies, the so-called Fanaroff-Riley (FR) Is (``low-power radio-loud AGN'' according to our nomenclature) overlap with that of radio-quiet AGN, which can generate some confusion and requires great care during the classification process, which needs to involve also the X-ray and far-IR bands (Section~\\ref{revised}). However, the two classes are physically distinct (see Sections~\\ref{2_lowrl} and \\ref{astro}), although the origin of radio emission in radio-quiet AGN is still not clear (but see Section~\\ref{rq_emission}).  Translating these high-level definitions into a classification scheme requires a wealth of multi-wavelength data, which were described in our previous papers. \\cite{kel08} (Paper I) presented the radio data of the VLA-CDFS sample, together with optical images and X-ray counterparts, while \\cite{mai08} (Paper II), discussed the optical and near IR counterparts to the observed radio sources and, based on rest-frame colors and the morphology of the host galaxies, found evidences for a change in the sub-millijansky radio source population below $\\approx 80$ microJy. \\cite{toz09} (Paper III), dealt with the X-ray properties, while \\cite{pad09} (Paper IV), discussed the source population.  This turned out to be made up of SFG and AGN at roughly equal levels, with the AGN including radio galaxies, mostly low-power (FR Is), and a significant ($\\sim 50\\%$) radio-quiet component. Paper IV made also clear that the ``standard'' definitions of radio-loudness, based on radio-to-optical flux density ratios, $R$, and radio powers, were insufficient to identify radio-quiet AGN when dealing with a sample, which included also star-forming and radio-galaxies, as both classes are or can be (respectively) characterized by low $R$ and radio powers as well. $R$, for example, is useful for quasar samples, where it can be assumed that the optical flux is related to the accretion disk, but obviously loses its meaning as an indicator of jet strength if both the radio and the optical band are dominated by jet emission, as might be the case in FR Is \\citep{chi99}.  Source classification in Paper IV was then based on radio, optical, and X-ray data, and was meant to provide a robust upper limit to the fraction of SFG at sub-millijansky levels. SFG candidates were selected based on their (low values of) $R$, (low) radio power, (non-elliptical or S0) optical morphology, and (low) X-ray power ($L_{\\rm x}$). The fact that X-ray upper limits above the AGN threshold ($10^{42}$ ergs s$^{-1}$) were also included was conservative in the sense that it maximized the number of SFG, as some of these sources could still be AGN. Furthermore, the selection of radio-quiet AGN candidates was only approximate, as it was based solely on $R$ and $L_{\\rm x}$ and suffered from uncertainties in the optical K-correction, possible contamination by radio galaxies, and the exclusion of X-ray upper limits (see Paper IV for details). In order to deal with the evolution and luminosity functions of the various classes of sources, we need to refine our classification. In particular, the CDFS field has been observed by {\\it Spitzer} and therefore near- (Section \\ref{nearIR}) and mid/far-IR (Section \\ref{farIR}) data are available for our sample.  The source classification used in this paper relies then on a combination of radio, IR, optical, and X-ray data (Section~\\ref{revised}).   Section~2 describes the updated classification of the VLA-CDFS sample, while Section~3 studies its evolution. Section~4 derives the LFs for various classes, while Section~5 discusses our results. Finally, Section~6 summarizes our conclusions. Throughout this paper spectral indices are defined by $S_{\\nu} \\propto \\nu^{-\\alpha}$ and the values $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.3$, and $\\Omega_{\\rm \\Lambda} = 0.7$ have been used. ",
        "conclusions": "We have used a deep, complete radio sample of 193 objects down to a 1.4 GHz flux density of $43~\\mu$Jy selected in the Chandra Deep Field South area to sharpen our understanding of the nature of sub-millijansky sources and to study for the first time their evolution and luminosity functions up to $z \\sim 5$. Our unique set of ancillary data, which includes far-IR, near-IR, and optical observations, redshift information, and X-ray detections or upper limits for a large fraction of our sources, has allowed us to develop an unprecedented classification scheme to categorize in a robust way faint radio sources in star-forming galaxies, radio-quiet, and radio-loud AGN. Our main results can be summarized as follows:  \\begin{enumerate} \\item Star-forming galaxies and AGN make up an equal part of the sub-millijansky sky down to $43~\\mu$Jy, with the former becoming the dominant population only below $\\approx 0.1$ mJy. Radio-quiet AGN are confirmed to be an important class of sub-millijansky sources, accounting for $\\sim 30\\%$ of the sample and $\\sim 60\\%$ of all AGN, and outnumbering radio-loud AGN at $\\la 0.1$ mJy. \\item The radio power of star-forming galaxies evolves as $(1+z)^{2.5 - 2.9}$ up to $z \\le 2.3$, their maximum redshift in our sample, in agreement with previous determinations in the radio and IR bands. Although evidence of a slowing down of the evolution at $z \\ga 1.3$ is not significant it cannot also be ruled out. The radio luminosity function of SFG can be parametrized as broken power-law $\\Phi(P) \\propto P^{-1.3}$ and $\\Phi(P) \\propto P^{-3.15}$ at the faint and bright end respectively, with a break at $P \\sim 7 \\times 10^{21}$ W~Hz$^{-1}$, which is also consistent with previous derivations. \\item AGN as a whole do not appear to evolve. However, once they are split into radio-quiet (energy budget dominated by thermal emission) and radio-loud (dominated by non-thermal, jet emission) sources the situation is very different, with the radio-quiet population evolving very significantly and similarly to star-forming galaxies and the radio-loud population displaying negative density evolution $\\Phi(z) \\propto (1+z)^{-1.8\\pm0.4}$. The luminosity function of radio-loud AGN, $\\Phi(P) \\propto P^{-1.5}$, is also much flatter than that of radio-quiet AGN, $\\Phi(P) \\propto P^{-2.6}$, derived here for the first time, which seems to be an extension of that of star-forming galaxies at the high power end. \\item There is a significant difference between the evolutionary properties of low-luminosity radio galaxies and radio powerful ($P \\ga 3 \\times 10^{24}$ W Hz$^{-1}$) AGN, as while the former do not evolve, the latter evolve negatively. This is exactly the opposite of what found in samples at higher flux densities ($\\ga 1$ Jy), where high-power sources exhibit a strong positive evolution. We interpret this difference as due to the fact that we are sampling the high-power radio-loud population up to much larger redshifts ($z \\ga 5$) and as a result of a redshift cutoff. \\item Our results suggest a very close relationship between star formation and radio emission in radio-quiet AGN, since their evolution and luminosity function are respectively indistinguishable from, and an extension of, that of star-forming galaxies. This is supported by the fact that radio-quiet AGN and star-forming galaxies appear both to follow the ``IR-radio relation'' but is in contradiction with the detection of compact, high brightness temperature cores in several (mostly local) radio-quiet AGN, similar to those observed in radio-loud ones. The co-existence of two components, one non-evolving and AGN-related, and one evolving and star-formation-related, and selection effects in the choice of radio-quiet VLBI targets, can reconcile these apparently discrepant observational data. \\item The surface density of radio-selected, radio-quiet AGN, $\\sim 520$ deg$^{-2}$, is already about 6 times larger than that of one of the deepest optically-selected quasar sample and only $\\sim 1/4$ that of unabsorbed X-ray selected AGN. This means that sub-millijansky radio surveys, given the appropriate ancillary multi-wavelength data, have the potential of detecting large numbers of radio-quiet AGN by-passing the problems of obscuration, which plague the optical and soft X-ray bands. The radio number counts of radio-quiet AGN are consistent with those in the hard X-ray band, which shows that the sources we are selecting are the same as the X-ray emitting radio-quiet AGN. \\item Sub-millijansky radio surveys wanting to study the evolution of faint, {\\it radio-loud} AGN, need to consider that a large fraction ($\\sim 60\\%$ down to $\\sim 50~\\mu$Jy) of the radio-selected AGN are actually of the {\\it radio-quiet} type and therefore need to be treated separately. This has not been done so far by other studies and can have a large impact on the study of faint ``radio-mode'' inefficient accretors.  \\end{enumerate}  We plan to expand on this work by using our deeper radio observations \\citep[][and in preparation]{mil08} and the recently released 4 Msec Chandra data \\citep{xue11}. This will provide us with a catalogue of $\\sim 900$ radio sources, with which we will be able to address the issues discussed in this paper with larger statistics."
    },
    "1107/1107.5179_arXiv.txt": {
        "abstract": "{} {In a previous paper we proposed a new model for the emission by amorphous astronomical dust grains, based on solid-state physics. The model uses a description of the disordered charge distribution (DCD) combined with the presence of two-level systems (TLS) defects in the amorphous solid composing the grains. The goal of this paper is to compare this new model to astronomical observations of different Galactic environments in the far-infrared/submillimeter, in order to derive a set of canonical model parameters to be used as a Galactic reference to be compared to in future Galactic and extragalactic studies.} {We compare the \\TLS model with existing astronomical data. We consider the average emission spectrum at high latitudes in our Galaxy as measured with FIRAS and WMAP, as well as the emission from Galactic compact sources observed with the Archeops balloon experiment, for which an inverse relationship between the dust temperature and the emissivity spectral index ($\\beta$) has been proven.} {We show that, unlike models previously proposed that often invoke two dust components at different temperatures, the \\TLS model successfully reproduces both the shape of the Galactic spectral energy distribution and its evolution with temperature as observed in the Archeops data. The best \\TLS model parameters indicate a charge coherence length of $\\simeq$ 13 nm and other model parameters in  broad agreement with expectations from laboratory studies of dust analogs.  We conclude that the millimeter excess emission, which is often attributed to the presence of very cold dust in the diffuse ISM, is very likely caused solely by TLS emission in disordered amorphous dust grains. We discuss the implications of the new model, in terms of mass determinations from millimeter continuum observations and the expected variations in the emissivity spectral index with wavelength and dust temperature. The implications for analyzing of the Herschel and Planck satellite data are discussed.} {} ",
        "introduction": "Early astronomical observations in the far-infrared (FIR) and submillimeter (submm) range have revealed the spectral energy distribution (SED) of dust emission towards the diffuse interstellar medium (ISM) to be consistent with a modified black-body function of the form \\begin{equation} \\label{eq:I} I_{\\nu}(\\lambda)=\\dustem\\,B_{\\nu}(\\lambda,T_d) \\times N_H \\end{equation} where $\\rm I_{\\nu}$ is the specific intensity or sky brightness (energy flux per unit wavelength and solid angle), $\\rm B_{\\nu}$ the Planck function, $\\lambda$ the wavelength, $\\rm T_d$ the dust equilibrium temperature, and $\\rm N_H$ the hydrogen column density. Those measurements were found to be consistent with a power-law emissivity function: \\begin{equation} \\label{eq_eps} \\dustem=\\epsilon(\\lambda_0) \\left( \\frac{\\lambda}{\\lambda_0} \\right )^{-\\beta} \\end{equation} where  $\\epsilon(\\lambda_0)$ is the emissivity at the reference wavelength $\\lambda_0$, and $\\beta$ is the emissivity spectral index, usually taken as $\\beta=2$ (the so-called quadratic law). The quality of the available data generally did not allow checking for variations in $\\beta$, which was therefore assumed to be constant.  The emission originates in dust radiating at thermal equilibrium (at the temperature $T_d$) with the surrounding radiation field, and as a consequence is related to the FIR and submm optical properties of the dust. Considering a constant gas/dust ratio $N_H/M_d$ (with $M_d$ the mass column density), and assuming a single dust temperature along the line of sight (LOS), the specific intensity can be written as \\begin{equation} I_{\\nu}(\\lambda)=\\kappa(\\lambda_0) \\left( \\frac{\\lambda}{\\lambda_0} \\right )^{-\\beta}B_{\\nu}(\\lambda,T_d)\\times M_d, \\end{equation} with \\begin{equation} \\kappa(\\lambda_0)=\\epsilon(\\lambda_0) \\frac{N_H}{M_d} \\end{equation} where $\\kappa(\\lambda_0)$ is the dust mass opacity at the reference wavelength $\\lambda_0$.  In solid state physics, within the framework of classical models, a constant and temperature-independent emissivity spectral index value $\\beta=2$ is justified by the asymptotic behavior of the absorption cross-section of solids at long wavelengths, far from the absorption resonances, such as the silicate absorption features at 10-20 $\\mic$. However, it was recognized early that the quadratic law may not hold for amorphous materials in the FIR/submm, since such materials are likely to exhibit low-energy transitions that happen in the FIR/submm \\citep[see][]{Andriesse74}. Although big grains (BGs) are known to be amorphous in nature in the diffuse ISM \\citep{Kemper04}, no need was identified to explicitly incorporate details about their internal structure in astrophysical models due to the lack of submm data. However, recent observations have proven that the actual FIR/submm dust SED is significantly more complicated than described by Eq. \\ref{eq_eps}.  Several recent studies have shown evidence of variations in the dust emissivity law with wavelength. That the overall Galactic SED is flatter than predicted by the quadratic law was clearly shown in the FIRAS data by \\citet{Reach95}. \\citet{Paradis09} studied the spectral dependence of dust emission in a collection of molecular clouds and their atomic surroundings over the spectral range covered by the DIRBE, Archeops, and WMAP experiments and showed that the observed emissivity law exhibits a break around 500 $\\mic$, with a steeper index ($\\beta$=2.4) in the FIR and a flatter index ($\\beta$=1.5) in the submm.  More recently a similar trend has been apparent in the Planck data of the Taurus molecular cloud \\citep{Abergel11}. Similar submm excess emission has also been observed towards external low-metallicity galaxies \\citet{Galliano05}, as well as towards the Small Magellanic Cloud in the Planck data \\citep{Bernard11}. The comparison between the spectral indices derived for the Galactic solar neighborhood ($\\beta=1.8$), the LMC ($\\beta=1.5$), and the SMC ($\\beta=1.2$) in the Planck data by \\citet{Planck11,Bernard11} indicates a general flattening that could be linked to the varying metallicity among these three environments. It therefore appears that the dust SED is usually flatter than predicted by the quadratic law and that the associated emissivity index $\\beta$ varies with wavelength and from object to object, possibly tracing evolution of the dust properties with metallicity.  In addition to the above, the spectral index of the dust emission as derived from a modified black-body fit to the data follows a systematic trend with dust temperature, the emissivity spectrum appearing flatter for regions with warmer dust temperatures and steeper for colder regions. \\begin{itemize}  \\item{\\citet{Dupac03} show variations in $\\beta$ in the FIR in various regions of the ISM, using the PRONAOS balloon-borne experiment data. They found variations in the index in the range 2.4 to 0.8 for $\\rm T_d$ between 11 and 80 K.} \\item{\\citet{Desert08} highlighted variations in $\\beta$ in the submm towards compact sources in Archeops balloon-borne experiment data. They found variations in the index in the range 4 to 1 for temperatures between 7 and 27 K.} \\item{\\citet{Veneziani10} show variations in $\\beta$ between 5 and 1 in the temperature range 7-20 K by analyzing 8 high Galactic latitudes clouds using BOOMERanG observations, combined with the IRAS, DIRBE, and WMAP data.} \\item{\\citet{Paradis10} found evidence of $\\rm T_d$-$\\beta$ variations using the combination of the new Herschel and IRAS data, in the inner regions of the Galactic plane. The $\\beta$ range is 2.6-1.8 for temperatures between 14 and 23 K.}  \\end{itemize} The origin of the flatness of the dust emission spectrum observed with FIRAS has been explored by \\citet{Reach95} and \\citet{Finkbeiner99}. The latter group propose a phenomenological description of the FIRAS data that invokes the coexistence of a warm ($T_d \\simeq16\\,K$) and a very cold dust ($T_d \\simeq 9\\,K$) component (hereafter FDS model). \\citet{Reach95} determined that the FIR and the submm emission were tightly correlated and therefore did not favor the very cold dust component.  Moreover, the two component description does not reproduce the observed variations in the emissivity spectral index with dust temperature. Indeed the spectral index associated to each of the components is set to a fixed value in the FDS model. For instance, in their best-fit model, $\\beta$ for the cold silicate-like component is set to 1.7,  which does not allow SEDs to be produced with a larger apparent emissivity index in the submm/mm domain (see Section \\ref{sec_beta_predictions}).  The two-level systems (\\TLS) model \\citep[][hereafter Paper I]{Meny07} has been developed to describe the FIR to mm continuum interstellar dust emission, taking the effect of the disordered internal structure of the amorphous dust grains into account. It is based on the solid-state physics model developed to interpret specific properties of the amorphous solids identified in laboratory data. As a consequence, it is expected to apply with a high degree of universality, and not to be sensitive to the exact chemical nature of the dust. The disordered charge distribution (DCD) part of the model describes the interaction between the electromagnetic wave and acoustic oscillations in the disordered charge of the amorphous material \\citep{Vinogradov60,Schlomann64}. This charge-disorder is observed on nanometer scale and is described here by a single charge correlation length. The TLS part of the model describes the interaction of the electromagnetic wave with a simple distribution of asymmetric double-well potential (ADWP), which results from disorder on atomic scale: a ADWP can be associated with two close configurations of atoms or group of atoms in the disordered structure \\citep{Phillips72, Phillips87, Anderson72}, where the corresponding atoms have two equilibrium positions and can transit from one to the other. Both DCD and TLS phenomena have been first applied by \\citet{Bosch77} to explain the observed temperature dependence of the absorption of some silica-based glasses in the FIR/mm.  In section 2, we present the data we consider for this analysis. In section 3, we briefly describe the model in terms of the disordered structure, as well as the main parameters of the model. In section 4, we explain the methods we used to compare the model to astrophysical data and derive standard parameters for the Milky Way (MW).  In section 5, we present model predictions for the standard parameters derived above, such as the evolution of the dust emission with temperature, the emissivity spectral index as a function of wavelength and temperature, and the predicted emissivity in Herschel and Planck photometric bands.  Sections 6 and 7 are devoted to the discussion and conclusions, respectively. ",
        "conclusions": "The theoretical \\TLS model described in a previous work (Paper I) and based on the amorphous structure of grains was compared to astrophysical data, such as the FIRAS/WMAP and Archeops data. The FIRAS/WMAP spectrum represents the diffuse interstellar medium in the stellar neighborhood, whereas the Archeops data characterize dust properties in a variety of compact sources, where a significant inverse relationship between the dust temperature and the emissivity spectral index had been proven. We performed a $\\chi^2$ minimization to determine standard values for the parameters of the \\TLS model selected to capture the spectral and temperature variations in the model. These free parameters are the dust temperature ($T_d$), the charge correlation length ($l_c$) that controls the wavelength where the inflection point between the two $\\beta=2$ and $\\beta=4$ ranges occurs, the intensity factor of the TLS processes ($A$) with respect to the DCD effect, and the intensity factor of the TLS/hopping process ($c_{\\Delta}$). Results indicate that emission in the submm/mm is dominated by the hopping relaxation. According to the model, the BG emission in the FIR/mm domain depends on wavelength and temperature, which is fundamental both for dust mass determination from FIR/submm measurements but also for component separation.  Using the best-fit parameters ($T_d=17.26$ K, $l_c$=13.4 nm, $A$=5.81, and $c_{\\Delta}$=475) allowing reproduction of both the emission from the diffuse medium and the compact sources, the model predicts significant $\\beta$ variations for temperatures between 5 and 100 K, with a maximum value of 2.6 at 2 mm. The \\TLS model is presently the only model able to predict $\\beta$ variations with temperature and wavelengths, as observed in both observational and laboratory data. The dust emissivity can be seriously underestimated if its variations with temperature and wavelength are not taken properly into account, generally inducing overestimates of the dust mass. Similarly, the comparison between temperatures derived from the \\TLS model ($\\rm T_{\\TLS}$) and those derived from a T$\\rm _d$-$\\beta$ fit ($\\rm T_{fit}$), between 100 and 500 $\\mic$, or 100 and 850 $\\mic$, are in good agreement for temperatures below 25 K. Above this value, the derived temperatures differ significantly from each other, in particular when emission far from the emission peak is considered in the fit. Indeed, a modified black-body fit is too simplistic to reproduce the mm emission excess, which increases with temperature. However, our study enables the possibility of deducing $\\rm T_{\\TLS}$ from a given $\\rm T_{fit}$. We also predicted dust emissivities in the IRAS 100 $\\mic$, Herschel, and Planck bands for $\\rm T_{\\TLS}$ between 5 and 100 K, which are useful for comparison with the Planck and Herschel data. For instance, the provided tables and figures  can be used to derive the charge correlation length and the intensity of the TLS processes from the Herschel and Planck data.  The \\TLS model is the first physically-motivated model able to reproduce both diffuse medium and compact sources, observed with FIRAS/WMAP and Archeops. In the future, the Planck data will be used to systematically compare the model prediction with submm data."
    },
    "1107/1107.2045_arXiv.txt": {
        "abstract": "{The near infrared (nIR)/optical counterparts of low mass X-ray binaries (\\lmxbs) are often observationally dim and reside in high source density fields which make their identification problematic; however, without such a counterpart identification we are unable to investigate many of the properties of \\lmxb\\ systems.} {Here, in the context of a larger identification campaign, we examine the fields of four \\lmxb\\ systems near the Galactic centre, in a bid to identify nIR/optical counterparts to the previously detected X-ray point sources.} {We obtain nIR/optical images of the fields with the ESO -- New Technology Telescope and apply standard photometric and astrometric calibrations; these data are supplemented by \\spitzer -\\glimpse\\ catalog data. } {On the basis of positional coincidence with the arcsecond accurate X-ray positions,  we identify unambiguous counterpart candidates for \\object{\\xtesixteen}, \\object{\\igrseventeen}, \\object{\\igrseventeenfive}\\ and \\object{\\gxnine}. } {We propose tentative nIR counterparts of four \\lmxbs\\  which require further investigation to confirm their associations to the X-ray sources. } ",
        "introduction": "\\label{section:intro}    Given the main sequence or red giant nature of the companion stars of low mass X-ray binaries (\\lmxbs), they are intrinsically dim in the  optical and near-infrared (nIR), and hence difficult to detect (see e.g., \\citealt {Charles2003:astro.ph.802} for a review of the optical and nIR properties of \\lmxbs). The intrinsic dimness of all these sources is further compounded by the high level of optical extinction in the direction of the Galactic plane \\citep{schlegel1998:ApJ500}, where many \\lmxbs\\ are found \\citep{Grimm2002:A&A391}; hence observing at nIR wavelengths, where this is less pronounced, may increase the chance of a detection. Due to their positions in the Galactic plane or centre, where the optical/nIR (OIR) source density is extremely high, accurate X-ray or radio positions are necessary to confidently identify the OIR counterparts of \\lmxbs\\ and reduce the probability of chance superpositions. The required (sub)arcsecond positions are now more readily available due to missions such as \\swift, XMM-{\\it Newton} and \\chandra, and when combined with the current generation of optical telescopes and instruments, increase the likelihood of counterpart discovery.    If the \\lmxb\\ is a persistent source or a transient source in outburst \\citep{vanParadijs1994:A&A.290} emission from the active accretion disk may be brighter than the companion star and more likely to be detected at OIR wavelengths; at nIR wavelengths it is even possible that a jet component will make a significant contribution to the emission (e.g., \\citealt{Russell2010:arXiv1001}). Broadband SED fitting of observed OIR counterparts can allow the various emission components (companion star and accretion disk/jet if present) to be disentangled from each other and, along with spectral lines, allows the spectral type -- and hence mass -- of the companion to be constrained. When combined with radial velocity information from X-ray or optical light curves this allows constraints to be placed on the mass of the compact object with which to confirm the nature of the compact object as derived from X-ray spectral or timing properties.  Furthermore the relationship between the OIR and X-ray flux can be used as a diagnostic with which to derive the compact object or companion star nature, and to identify the state of the system at the time of observations (e.g., \\citealt{russell2006:MNRAS371}).     \\begin{table*} \\centering \\caption{X-ray and nIR positions for the \\lmxbs\\ in our study} \\label{table:positions} \\begin{tabular}{l l l l l l l l} \\hline\\hline & X-ray &  &  & & nIR &  &  \\\\ Source & RA & Declination & Err (\\arcsec) & & RA & Declination & Err (\\arcsec) \\\\ \\hline \\xtesixteen$^1$ & 16:37:02.67  & $-$49:51:40.6 & 1.8  & & 16:37:02.66  & $-$49:51:40.0 & 0.16 \\\\ \\igrseventeen$^2$ & 17:37:58.81 &  $-$37:46:19.6 & 1.4 & &  17:37:58.77 &  $-$37:46:19.9 & 0.32 \\\\ \\igrseventeenfive$^3$ & 17:58:29.85 &  $-$30:57:01.6 & 0.6 & &  17:58:29.83 &  $-$30:57:01.6 & 0.16 \\\\ \\gxnine$^3$ & 18:01:32.15 &  $-$20:31:46.1 & 0.6 & &  18:01:32.15 &  $-$20:31:46.2 & 0.16 \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[]  Reference to X-ray positions: $^1$ \\cite{Starling2008ATel.1704}, $^2$ \\cite{Krimm2008:ATel.1714}, $^3$ M. Revnivtsev (private communication) \\end{list} \\end{table*}       Here we examine the fields of four \\lmxb\\ systems near the Galactic centre, in a bid to identify nIR/optical counterparts to the previously detected X-ray point sources (Table\\,\\ref{table:positions}). For \\xtesixteen, \\igrseventeen, \\igrseventeenfive\\ and \\gxnine\\ we identify unambiguous counterpart candidates on the basis of positional coincidence with the arcsecond accurate X-ray positions. We have also observed a further three \\lmxb\\ systems (\\object{\\onea}, \\object{\\slxoo}\\ and  \\object{\\igrseventeenofive})$^*$ as part of our campaign to identify the nIR counterparts of high energy sources near the Galactic centre. Though we do not detail these observations, they confirm the results of \\cite{Zolotukhin:2011MNRAS.411} and have been used to verify our reduction methods. In section \\ref{section:observations} we introduce our observations and reduction method while in section \\ref{section:results}, after briefly introducing each source,  we detail the results of those observations.  We summarise our findings in section \\ref{section:conclusions}.  Throughout, positions (J2000) are given with 90\\% confidence while all others values, including magnitudes, are given with $1\\sigma$ confidence. ",
        "conclusions": "\\label{section:conclusions}  To approximate the probability of a chance superposition of the 90\\% X-ray positions with random sources in the respective fields, we calculate $P \\approx \\rho_N \\times A_{Err}$; where $\\rho_N$ is the surface area number density of observed sources down to the limiting magnitude and $A_{Err}$ is the area of the X-ray positional error. The probability of chance coincidence for \\xtesixteen\\ and \\igrseventeen\\ are  quite high at 86\\% and 28\\% respectively, due to a large number of sources in the deep fields but more significantly, the radius of the X-ray positional errors. The importance of sub-arcsecond positions from e.g. \\chandra\\ is further illustrated by the significantly lower probabilities of \\igrseventeenfive\\ and \\gxnine\\ (8\\% and 7\\%) which have comparable number densities to the previous two fields but much more accurate X-ray positions. The probability of chance superposition could be reduced further if we limited our search to sources of specific magnitude ranges but due to a lack of information on these sources, we make no such assumption.   On the basis of positional coincidence with the arcsecond accurate X-ray positions, we suggest unambiguous (in so far as there is only one per source) counterpart candidates for four \\lmxbs\\ near the Galactic centre. We caution that given the probability of chance superpositions in these crowded fields, the associations are tentative and require further investigations of e.g., spectral type, variability, distance, to confirm the associations with the X-ray sources. The spectral shapes of the counterpart candidates are not particularly constraining but, in general, could be consistent with stellar-like black body emission in the optical/nIR but with an excess at the longer, \\spitzer, wavelengths, as for \\xtesixteen. The excess could be due to warm circumbinary dust, likely related to the formation of the compact object (e.g., \\citealt{Rahoui2010:ApJ.715}) or to non-thermal, synchrotron jet emission (e.g., \\citealt{Gelino2010:ApJ.718}). Further, possibly simultaneous, photometric and spectroscopic observations are required to disentangle the different emission components and identify the spectral type of the companion star, if it is  detected over the other possible emission mechanisms."
    },
    "1107/1107.5193.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary { } % aims heading (mandatory) {Interstellar dust grains, because of their catalytic properties, are crucial to the formation of \\hm, the most abundant molecule in the Universe. The formation of molecular hydrogen strongly depends on the ability of H atoms to stick on dust grains. In this study we determine the sticking coefficient of H atoms \\rm{chemisorbed}\\rm\\ on graphitic surfaces, and estimate its impact on the formation of \\hm. } % methods heading (mandatory) {The sticking probability of H atoms chemisorbed onto graphitic surfaces is obtained using  a mixed classical-quantum dynamics method. In this, the H atom is treated quantum-mechanically and the vibrational modes of the surface are treated classically. The implications of sticking for the formation of \\hm\\  are addressed by using \\rm{Kinetic}\\rm\\ Monte Carlo simulations that follow how atoms stick, move and associate with each other on dust surfaces of different temperature. } % results heading (mandatory) {\\rm{In our model}\\rm, molecular hydrogen forms very efficiently for dust temperatures lower than 15~K through the involvement of physisorbed H atoms. At dust temperatures higher than 15~K and gas temperatures lower than 2000~K, \\hm\\ formation differs strongly if the H atoms coming from the gas phase have to cross a square barrier (usually considered in previous studies) or  a barrier obtained by DFT calculations to become chemisorbed. The product of sticking times efficiency can be increased by many orders of magnitude when realistic barriers are considered. If graphite phonons are taken into account in the dynamics calculations, then H atoms stick better on the surface at high energies, but the overall \\hm\\ formation efficiency is only slightly affected. Our results suggest that \\hm\\ formation can proceed efficiently in photon dominated regions, X-ray dominated regions, hot cores and in the early Universe when the first dust is available. } % conclusions heading (optional), leave it empty if necessary { } ",
        "introduction": "Interstellar dust grains play a very important role in the chemistry of the interstellar medium (ISM). Because of inefficient gas phase paths to form \\hm, dust grains are considered as the favored habitat to form \\hm\\ molecules (\\citealt{gould1963}). In the past decades, a plethora of laboratory experiments and theoretical models have been developed in order to understand how the most abundant molecule of the Universe forms. The sticking of H atoms on surfaces has received considerable attention since this mechanism governs the formation of H$_2$, but also other molecules that contain H atoms. The sticking of H atoms on dust grains can also be an important mechanism to cool interstellar gas (\\citealt{spaans2000}). As H atoms arrive on the dust, they can be weakly bound (physisorbed) or strongly bound (chemisorbed) to the surface. \\rm{The sticking of H in physisorption state has been highlighted by several experiments on different type of surfaces (\\citealt{pirronello1997a}, \\citeyear{pirronello1999}, \\citeyear{pirronello2000},\\citealt{perry2003}). Experiments performed on surfaces at higher temperatures revealed how H atoms stick in chemisorbed state (\\citealt{zecho2002}, \\citealt{hornekaer2006}, \\citealt{mennella2006}). In addition, theoretical studies confirmed that the sticking of H atoms  varies strongly with gas and dust temperatures (\\citealt{leitch1985}, \\citealt{buch1989}, \\citealt{sha2005}, \\citealt{medina2008}, \\citealt{morissetbis2008}, \\citealt{cuppen2010}, \\citealt{morisset2010}})\\rm. To study the sticking of an atom on a surface, it is necessary to take into account the vibrational modes (phonons) of the surface (\\citealt{burke1983}, \\citealt{buch1989}, \\citealt{morissetbis2008}, \\citealt{morisset2010} ).  In the ISM, dust grains are mainly carbonaceous particles or silicates, with various sizes, and a large fraction of the available surface for chemistry is in the form of very small grains or polyaromatic hydrocarbons (PAHs; \\citealt{Weingartner2001}). PAHs have similar characteristics as graphite surfaces, as shown by different calculations (\\citealt{jeloaica1999}, \\citealt{sha2002}, \\citealt{ferro2008}). \\cite{jeloaica1999} studied the interaction between H atoms and a coronene,  whereas \\cite{ferro2008}  and \\cite{sha2002}  determined the H-graphite interaction. All these theoretical works obtained identical  characteristics: the H atom coming from the gas phase can physisorb (43~meV) or chemisorb (0.76~eV) on the graphite or PAH. \\rm{Several experimental studies showed that H atoms can physisorb on carbonaceous (\\citealt{pirronello1997a}, \\citeyear{pirronello1999}, \\citeyear{pirronello2000}), silicates (\\citealt{pirronello1997b})  and graphitic (\\citealt{perry2003}) surfaces, and  can also chemisorb on graphite surfaces  (\\citealt{zecho2002}, \\citealt{hornekaer2006}). These experiments revealed that adsorption of H atoms (physisorption and chemisorption) is of key importance to form H$_2$ for a wide range of dust and gas temperatures, as observed in the ISM.}\\rm\\ The H atom can chemisorb only on top of a C atom and a surface reconstruction is observed: the C atom goes out of the surface plane. This phenomenon is called puckering, and creates a barrier against chemisorption of 0.2~eV (\\citealt{jeloaica1999}, \\citealt{sha2002}, \\citealt{ferro2003}, \\citealt{hornekaer2006}). This barrier is important and does not allow an efficient sticking mechanism for atoms coming in at low energies. Therefore, sticking through chemisorption strongly depends on the energy of the incoming atoms. Also, the sticking is sensitive to the temperature of the dust since the surface excitation modes play a role in redistributing the excess energy of an atom. \\rm      The \\hm\\ formation occurs through Eley-Rideal (ER) or Langmuir-Hinshelwood (LH) mechanisms. In the ER mechanism, an H  atom coming from the gas phase collides with an H atom which is initially adsorbed (chemisorbed) on the graphite surface. Theoretical (\\citealt{rougeau2006}, \\citealt{ferro2008}) and experimental (\\citealt{hornekaer2006}) studies show that H atoms are chemisorbed on the surface in dimer  configurations, and that two hydrogen dimers configurations are stable: the ortho-dimer and the para dimer (Figure 3, right panel). Density Functional Theory (DFT) calculations show that while the chemisorption of one H atom is associated with an important barrier, the formation of the para dimer is barrier-less and that the formation of the ortho dimer has a reduced barrier (\\cite{rougeau2006}). The formation of \\hm\\ that involves chemisorbed atoms through the ER mechanism has been studied by Density Functional Theory (DFT) (\\citealt{jeloaica1999}, \\citealt{sha2002}, \\citealt{ferro2003}) and dynamics calculations (\\citealt{rutigliano2001}, \\citealt{ree2002}, \\citealt{sha2002},  \\citealt{morisset2003}, \\citeyear{morisset2004b}, \\citealt{martinazzo2006chem}). Based on these calculations, different mechanisms have been proposed to contribute to the \\hm\\ formation through the ER mechanisms: the direct ER that involves isolated H atoms (monomers,  \\citealt{sha2002}, \\citealt{morisset2003}, \\citeyear{morisset2004b}, \\citealt{martinazzo2006chem}), barrier-less formation of \\hm\\ involving one H atom in a para-dimer configuration (\\citealt{bachellerie2007}) and formation by diffusing H atoms in physisorbed states  (\\citealt{bonfanti2007}). In the LH mechanism, the two H atoms are adsorbed (physisorbed) on the graphite surface, diffuse on the surface and collide to desorb in a hydrogen molecule.  This mechanism has also been studied (\\citealt{morisset2004},  \\citeyear{morisset2005},  \\citealt{martinazzo2006phys}). The different mechanisms to form \\hm\\ on graphite surfaces are the following:\\\\ M1) LH mechanism: Two physisorbed H atoms encounter each other on the surface; \\\\ M2) direct Eley-Rideal mechanism involving \\rm{H atom chemisorbed in a monomer (only one H atom chemisorbed on the cycle, fig. \\ref{benzene}, right panel)}\\rm;\\\\ M3) direct Eley-Rideal mechanism involving an H atom chemisorbed in a  dimer \\rm{(H atom  chemisorbed in dimer position on the first cycle, fig. \\ref{benzene}, right panel)\\rm;\\\\ M4) fast diffusion of physisorbed H atoms that enter chemisorbed sites occupied by H atoms (monomer);\\\\ M5) fast diffusion of physisorbed H atoms that enter chemisorbed sites occupied by H atoms (dimer);\\\\ M6) direct Eley-Rideal mechanism involving an H atom in meta, 5 and S (second cycle) positions (\\citealt{dumont2008}, see figure 3 right panel);\\\\ M7) ER mechanism by fast diffusing H atoms in the physisorption state (\\citealt{bonfanti2007}).\\\\  \\rm{Recent experimental studies (\\citealt{lemaire2010}, \\citealt{islam2007}) show that the ro-vibrationnal states of the \\hm molecule formed on the surface can be detected (silicates and graphite). The ro-vibrationnal distribution of the \\hm\\ allows to classify the LH and the ER mechanisms to form the molecule as function of the dust temperature. Experimentally, it is not possible to identify the M1 to M7 mechanisms proposed to form \\hm. The objective of our work is to classify the predominant mechanisms to form \\hm\\ as a function of gas and dust temperature.}\\rm\\ In each mechanism proposed, the sticking of the H atoms in physisorption sites and chemisorption sites (figure 3 left panel) is crucial to allow the \\hm\\ formation. In this study, we concentrate on understanding how the first H atom chemisorbs on the grain with a direct ER mechanism and how the sticking probability impacts the \\hm\\ formation. During the collision between an atom and a surface, the collisional energy is distributed between the vibrational excitation of the newly formed bond and the vibrational modes of the surface.Therefore, to calculate the sticking of atoms on a surface in a realistic way, one has to consider the lattice dynamics. For this purpose, we study the interaction of H atoms with graphitic surfaces, which are surfaces having similar properties as PAH surfaces. Our attention focusses particularly on these surfaces because PAHs represent a large fraction of the total surface area of dust grains (up to 50$\\%$, \\citealt{Weingartner2001}).  In the first section, the time dependent dynamics method is presented to study the sticking of H atoms in chemisorbed sites on the graphite surface in the collinear geometry. In the subsequent section, the different mechanisms for the formation of \\hm\\ are implemented into the Kinetic Monte Carlo (KMC) simulations. The sticking probability obtained by the dynamics method is included in the KMC code.  In the last section, we assess how this sticking changes the processes involved in the formation of H$_2$ \\rm{and identify the predominant mechanisms to form \\hm\\ molecules as function of the gas and dust temperature.}\\rm ",
        "conclusions": "In this work, we have studied the influence of the sticking probability on the formation of  H$_2$ \\rm{using Kinetic Monte Carlo simulations. In our model, mechanisms M1 to M5 (described in the introduction) for the \\hm\\ formation are included. We have determined the predominant mechanisms to form \\hm\\ as function of gas and dust temperatures.  First, we assumed the sticking probability of H atoms to be equal to the transmission coefficient to cross the barrier against chemisorption. To address the importance of the shape of the barrier, we consider a square barrier, and a barrier obtained by DFT calculations. For the latter, wavepacket propagation calculations along the coordinate between the H atom and the graphite surface (collinear approach) are performed to determine the transmission probabilities of the chemisorption barrier. Our results show that if a square barrier is considered, the transmission probabilities are underestimated by many orders of magnitude at low gas temperatures (T$_{gas}$$\\le$2000K). The impact of the form of the barrier is visible on the results of the KMC calculations simulating the formation of \\hm. Our results show that the shape of the chemisorption barrier (square or obtained by DFT) has a strong impact on the \\hm\\ formation for T$_{dust}$ $\\ge$ 15~K. This impact increases as gas temperatures decrease. In conclusion, the formation of \\hm\\ calculated with barriers obtained by DFT is much more efficient than with square barriers. In these temperatures domain, \\hm\\ formation involves chemisorbed atoms. For T$_{dust}$$\\le$15~K, \\hm\\ is formed by a physisorption mechanism. In this case, the form of the chemisorption barrier is not important.  Second, we investigated the sticking of H atoms in chemisorbed sites taking into account the vibrational modes of the surface. For this we used a mixed classical-quantum dynamics method where the motion of the H atom is treated quantum-mechanically, while the vibrational modes of the surface are treated classically (\\citealt{morisset2008}). The sticking probability which take into account the vibrational modes of  the surface can be compared to the transmission probability of the barrier. The sticking probability is higher by a factor of 5 for T$_{gas}$ $\\sim$1500K, but is lower by a factor 2 (for T$_{dust}$=10~K) to 5 (for T$_{dust}$=125~K) around T$_{gas}$ $\\sim$1000K. This calculation show that the energy exchange between the atom and the bath of phonon is higher than transmission probability of the chemisorption barrier. These sticking probabilities are calculated as function of T$_{gas}$, for T$_{dust}$=10~K and T$_{dust}$=125~K. We implement these sticking coefficients S$_{chem}$ in our KMC  simulations, and show that for  T$_{dust}$=125~K, \\hm\\  is formed through a chemisorption mechanism ( involving dimers for T$_{gas}$$\\le$ 1000~K and monomers for  T$_{gas}$$\\ge$ 1000~K ). For T$_{dust}$=10~K, on the other hand, \\hm\\ formation involves physisorbed atoms if the sticking in physisorption is high. In this case the efficiency to form \\hm\\ depends of the sticking probabilities of H atom in a physisorption site.  \\rm %We investigated the effect of the sticking in chemisorbed sites on the formation of \\hm. It appears that at  dust temperatures higher than 15~K,  chemisorbed H atoms are involved in the formation of \\hm, and the product $S\\times \\epsilon$ depends on $S_{chem}$, the sticking in chemisorption. % We considered different sticking mechanisms in chemisorbed sites, $S_{chem}$, (1) obtained for a square barrier, (2)  obtained for a barrier determined by DFT calculations and (3) the same as (2) taking into account the vibrational modes of the surface. The product $S\\times \\epsilon$ obtained with a square barrier (1) is many orders of magnitudes lower than the one obtained with a realistic barrier (2), for low gas temperatures.  Once the vibrational modes are taken into account in the sticking (3), then the \\hm\\ formation efficiency is increased by up to a factor of 2 for T$_{gas} >$ 400~K.   The sticking in physisorption influences the formation of H$_2$ at low dust temperatures ($\\le$15~K). In this study, we considered 2 possible sticking coefficients in physisorption S$_{phys}$. 1) The sticking  derived in \\cite{cuppen2010}, which is valid for metals and is close to unity for low gas and dust temperatures. 2) The sticking derived by \\cite{buch1989},  which shows that the probability of sticking for H atoms on graphite is on the order of a few percent (\\citealt{lepetit2011}, \\citealt{medina2008}). The influence of this S$_{phys}$ on the product $S\\times \\epsilon$ is very important at low dust temperatures. If the sticking is taken as in 1), then \\hm\\ is formed through LH mechanisms with physisorbed atoms, and $S\\times \\epsilon$ is very high $\\sim$ S$_{phys}$. On the other hand, if the sticking in physisorption is chosen to be as in 2), the formation of \\hm\\ involving physisorbed atoms is not efficient, and the formation of \\hm\\ is insured by chemisorbed H atoms. In this case, the product $S\\times \\epsilon$ depends on S$_{chem} \\times \\epsilon_{chem}$. \\rm{Our KMC simulations show that  once the vibrational modes of the H-graphite system are taken into account in the sticking, the efficiency to form \\hm\\ is increased by a factor 2. This behavior is not surprising since the sticking probability obtained by dynamics calculations is higher than the value of the transmission probability of the chemisorption barrier. }\\rm  \\rm{In the ISM, the \\hm\\ formed are ro-vibrationally excited.  The excitation of the molecule is due to the energy transfer of the kinetic energy towards the substrate, the translational energy and the excitation of the molecule. The ro-vibrational excitation of the molecule allows the determination of the mechanisms to form \\hm\\ (\\citealt{lemaire2010}, \\citealt{islam2007}). Different mechanisms have been studied theoretically by several groups (\\citealt{rutigliano2001}, \\citealt{ree2002}, \\citealt{sha2002},\\citealt{morisset2003}, \\citealt{morisset2004}, \\citealt{morisset2005}, \\citealt{martinazzo2006phys},\\citealt{bachellerie2007}). In our KMC simulations it is not possible to extract the energy of the \\hm\\ formed. These simulations only allow us to follow the number of molecules formed through different mechanisms (M1 to M5 described in the introduction). In this study, we have determined the predominant mechanism to form \\hm\\ as a function of the dust and gas temperatures.}\\rm   Recently, \\cite{cuppen2010} performed similar work on the formation of \\hm\\ and the influence of the sticking on the product  $S\\times \\epsilon$. These authors concentrated on the sticking in physisorbed sites (and we used their results in this work), but assumed that the atoms from the gas phase had to thermally overcome the barrier against chemisorption to become chemisorbed. \\rm{In this sense, the H atoms coming from the gas phase do not stick in chemisorbed sites at low T$_{gas}$ (S$_{chem}$= 0.03, 0.0001 and 10$^{-8}$ for T$_{gas}$= 500, 200 and 100~K respectively). Therefore, our results are similar to \\cite{cuppen2010}  for cold dust, when physisorbed atoms are involved to form \\hm. However, on warm dust grains ($\\sim$ 100K), chemisorbed atoms are needed to form \\hm\\ and the sticking of H atoms in chemisorbed sites sets the formation of \\hm. In this case, our results diverge from \\cite{cuppen2010} for gas temperatures lower than 500~K, since  we consider that H atoms from the gas phase can tunnel through the important barrier against chemisorption}. This process makes the sticking of H atoms in chemisorbed sites still efficient at low T$_{gas}$. Once these chemisorbed sites are populated, \\hm\\ can form efficiently through barrier-less routes involving the creation of dimers. Therefore, our results show that the formation of \\hm\\ is efficient also for intermediate gas and dust temperatures (T$_{dust}$>15K and 1000~K>T$_{gas}$>100~K).\\rm  Our results suggest that the formation of H$_2$ remains efficient in regions where gas and dust are warm. \\rm{This is true if the time for H atoms to enter chemisorbed sites is shorter than other routes to form \\hm\\ (i.e. dense and warm medium). Therefore, the grain surface route proposed in this study should be compared to other routes to form \\hm. }\\rm\\ Environments where grain surface chemistry would dominate to form \\hm\\ include photon dominated regions (PDRs, \\citealt{hollenbach1999}), X-ray dominated regions (XDRs, \\citealt{meijerink2005}), hot cores (\\citealt{caselli1993}), and the early universe (\\citealt{cazaux2004b}, \\citeyear{cazaux2009}). Particularly PDRs and XDRs enjoy regions of upto $10^{22}$ and $10^{24}$ cm$^{-2}$, respectively, where rapid photo-dissociation and chemical removal of H$_2$ requires it to be formed efficiently on dust grains in order to drive ion-molecule chemistry and H$_2$ line emission (\\citealt{habart2004}). In PDRs and XDRs gas temperatures can be as high as $10^3$ K (\\citealt{meijerink2007}), values for which the traditional sticking coefficient of \\cite{hollenbach1979} decreases rapidly. At early cosmic epochs, for redshifts larger than 10, the first grain surfaces are expected to be warm, at least 30 K, due to the strong cosmic microwave background. Furthermore, the low abundances of metals like carbon and oxygen cause gas temperatures to be $10^2-10^{3.5}$ K and to be set mostly by H$_2$ and HD cooling (\\citealt{glover2008}). Our results indicate that chemisorption effects allow dust grains to act as catalysts for H$_2$ formation under such hostile primordial conditions.   \\subsection*"
    },
    "1107/1107.5515_arXiv.txt": {
        "abstract": "This is the first in a series of papers studying the variable stars in Large Magellanic Cloud globular clusters.  The primary goal of this series is to better understand how the RR Lyrae stars in Oosterhoff-intermediate systems compare to those in Oosterhoff I/II systems.  In this paper we present the results of our new time-series $BV$ photometric study of NGC 1466.  A total of $62$ variables were identified in the cluster, of which $16$ are new discoveries.  The variables include $30$ RRab stars, $11$ RRc's, $8$ RRd's, $1$ candidate RR Lyrae, $2$ long-period variables, $1$ potential anomalous Cepheid, and $9$ variables of undetermined classification.  We present photometric parameters for these variables.  For the RR Lyrae stars physical properties derived from Fourier analysis of their light curves are presented.  The RR Lyrae stars were used to determine a reddening-corrected distance modulus of $(m-M)_{0} = 18.43\\pm 0.15$.  We discuss several different indicators of Oosterhoff type and find NGC 1466 to be an Oosterhoff-intermediate object. ",
        "introduction": "Studies of RR Lyrae stars in the Milky Way globular clusters reveal what is traditionally known as the Oosterhoff dichotomy; the tendency for these clusters to be either Oosterhoff I (Oo-I) or Oosterhoff II (Oo-II) objects with a relatively clear zone of avoidance between these two groups.  The left panel of Figure 5 in \\citet{mc09} strikingly illustrates the Oosterhoff dichotomy in a plot of the average period of the RR Lyrae stars of Bailey type ab (RRab) versus cluster metallicity for Milky Way globular clusters.  When one looks at the nearby dwarf galaxies and their globular clusters, one does not see an Oosterhoff dichotomy as these objects fall not only into the Oo-I/II groups but also into the gap between those two groups.  In fact the distribution of these extragalactic objects seems to peak in the zone of avoidance \\citep{mc09}.  The existence of Oosterhoff-intermediate objects (Oo-int), objects that fall into the zone of avoidance, in the nearby dwarf galaxies poses a significant challenge to the hierarchical merger model for the formation of the Milky Way halo.  The lack of Oo-int clusters in the halo means that the objects that were accreted to form the halo could not have had properties entirely like the present-day nearby dwarf galaxies.  This is the first of a series of papers focusing on variable stars in globular clusters in the Large Magellanic Cloud (LMC).  The goal of this series is to better understand how the RR Lyrae stars in Oo-int systems compare to those in Oo-I/II systems.  The LMC is an ideal place for this study because of the large number of Oo-int clusters that it contains \\citep{bo94,mc09}; of the twelve LMC globular clusters that contain at least $5$ RRab variables, five of these clusters are Oo-int (NGC 1466, NGC 1853, NGC 2019, NGC 2210, and NGC 2257).  In this paper we identify, classify, and discuss the variable stars found in the LMC globular cluster NGC 1466.  Subsequent papers will focus on the variable stars in other LMC globular clusters.  NGC 1466 is an old globular cluster that is located relatively far from the center of the LMC.  It has a metal abundance of [Fe/H] $\\approx -1.60$, see discussion in section \\ref{sec:abprop}, and is not very reddened, E(B-V) = $0.09\\pm0.02$ \\citep{wa92b}.  \\citet{jj99} used color-magnitude diagrams obtained with the Hubble Space Telescope to determine that the age of NGC 1466 is within $1$ Gyr of the ages of the Milky Way globular clusters M3 (NGC 5272) and M92 (NGC 6341).  NGC 1466 features a well-populated horizontal branch that extends through the instability strip (see Figure 5 in Walker 1992); thus it is expected to contain a significant number of RR Lyrae stars. RR Lyrae stars were first found in NGC 1466 by \\citet{tw53}, and additional RR Lyrae were discovered by \\citet{we71}.  The most recent study of variables in NGC 1466 was conducted by \\citet{wa92b}, who found $42$ RR Lyrae stars; $25$ RRab and $17$ RRc stars.  Due to the density of unresolved stars in the cluster core, Walker did not perform photometry of stars within a radius of $13$ arcsec from the cluster center, suggesting that there are probably additional RR Lyrae stars that could not be detected in his study.  The core radius of NGC 1466 is $10.7\\pm0.4$ arcsec, as measured by \\citet{mg03} using images from the Hubble Space Telescope, which is entirely within the region of the cluster where Walker was unable to perform photometry.  Advances in instrument resolution and image-subtraction techniques now make it easier for us to search for variable stars in more crowded regions.  In this paper we present the results of a new search for variable stars in NGC 1466 and use the RR Lyrae in NGC 1466 to confirm the distance modulus and Oosterhoff classification of the cluster.  We also present Fourier fits to the light curves and derive physical properties for the stars from these fits. ",
        "conclusions": "A photometric study of the LMC globular cluster NGC 1466 was conducted in order to locate variable stars, with a specific interest in the cluster's RR Lyrae population.  $49$ RR Lyrae stars and $1$ candidate RR Lyrae were found.  We discovered $9$ RR Lyrae stars which were not found by \\citet{wa92b}, previously the most extensive search for RR Lyrae in NGC 1466.  Of our $49$ RR Lyrae stars, $8$ are RRd stars, making this the first time double-mode RR Lyrae stars were found in this cluster.  A candidate anomalous Cepheid variable was found in NGC 1466.  ACs are rare in globular clusters, having previously been found only in NGC 1786, NGC 5466, and $\\omega$ Cen \\citep{ne94,cc01}.  Interestingly, these three systems are OoII, though this of course may be just a coincidence, particularly since it is believed that ACs in globular clusters should be the progeny of blue straggler stars (e.g., Sills et al. 2009 and references therein), which in principle should not be related to the Oosterhoff dichotomy.  We performed Fourier analysis on the RR Lyrae light curves and used the resulting Fourier parameters to calculate physical properties for the stars which will be compared to the properties of RR Lyrae stars in other clusters in a future paper in this series.  We obtained a reddening-corrected distance modulus of $(m-M)_{0}=18.43\\pm0.17$ from the RRab and RRc stars.  The average periods for the RRab and RRc stars and the ratio of number of RRc to total number of RR Lyrae all suggest an Oosterhoff-intermediate classification for NGC 1466.  The position of the RRab stars in the Bailey diagram provides for a less clear Oosterhoff determination."
    },
    "1107/1107.4039_arXiv.txt": {
        "abstract": "Stellar rotation periods measured from single-age populations are critical for investigating how stellar angular momentum content evolves over time, how that evolution depends on mass, and how rotation influences the stellar dynamo and the magnetically heated chromosphere and corona. We report rotation periods for 40 late-K to mid-M stars members of the nearby, rich, intermediate-age ($\\sim$600 Myr) open cluster Praesepe. These rotation periods were derived from $\\sim$200 observations taken by the Palomar Transient Factory of four cluster fields from 2010 February to May. Our measurements indicate that Praesepe's mass-period relation transitions from a well-defined singular relation to a more scattered distribution of both fast and slow rotators at $\\sim$0.6 $\\Msun$. The location of this transition is broadly consistent with expectations based on observations of younger clusters and the assumption that stellar-spin down is the dominant mechanism influencing angular momentum evolution at 600 Myr. However, a comparison to data recently published for the Hyades, assumed to be coeval to Praesepe, indicates that the divergence from a singular mass-period relation occurs at different characteristic masses, strengthening the finding that Praesepe is the younger of the two clusters. We also use previously published relations describing the evolution of rotation periods as a function of color and mass to evolve the sample of Praesepe periods in time. Comparing the resulting predictions to periods measured in M35 and NGC 2516 ($\\sim$150 Myr) and for kinematically selected young and old field star populations suggests that stellar spin-down may progress more slowly than described by these relations. ",
        "introduction": "In a seminal paper, Andrew Skumanich (1972) showed that stellar rotation decreases over time such that $v_{rot} \\propto t^{-0.5}$, as does chromospheric activity, a proxy for magnetic field strength. This relationship between age, rotation, and activity has been a cornerstone of stellar evolution work over the past 40 years, and has generated almost as many questions as applications. Angular momentum loss due to stellar winds is generally thought to be responsible for the Skumanich law, but the exact dependence of $v_{rot}$ on age is not entirely understood, and relies on the assumed stellar magnetic field geometry and degree of core-envelope coupling \\citep{kawaler1988, krish1997}. Furthermore, later-type, fully convective stars appear to have longer active lifetimes than their early-type brethren \\citep[e.g.,][]{andy08}, indicating that they are capable of generating significant magnetic fields even in the absence of a standard solar-type dynamo \\citep{Browning2008}. The lack of a comprehensive theoretical understanding of the age-rotation-activity relation has not prevented the development and use of gyrochronology, however, which is used to determine the ages of field stars based on a presumed age-rotation relation \\citep[e.g.,][]{barnes2007, mamajek2008, collier2009, barnes2010}, nor of empirical age-activity relations, which do not always find activity decaying quite as simply as predicted by the Skumanich law \\citep[e.g.,][]{feigelson04,pace2004,giampapa2006}.  Mapping out the dependence of stellar rotation and activity on age requires the study of stars ranging in both mass and age. Statistical constraints on the age-rotation-activity relation can be derived from Galactic field stars \\citep[e.g.,][]{feigelson04, Covey2008, irwin2011}, but the homogeneous, coeval populations in open clusters provide an ideal environment for studying time-dependent stellar properties. Ideally, rotation periods for large numbers of cluster members could be measured directly from modulations in these stars' light-curves due to the presence of star spots. There are relatively few nearby open clusters, however, and fewer still have had the high quality photometric data needed to characterize their members' rotation in this manner --- in part because of the sheer difficulty involved in systematically monitoring a large number of stars over several months or more. Studies like that of \\cite{skumanich72} relied instead on measurements of the rotational Doppler broadening of spectral lines, a technique that has the advantage of needing only one observation. Translating the resulting $v_{rot}$ sin $i$ measurements into $P_{rot}$ involves making assumptions about stellar radii and inclinations, however, neither of which are well constrained.\\footnote{A further limitation of the Doppler broadening technique is that it is sensitive only to stars rotating faster than some threshold set by the spectral resolution.}  Because of these challenges, our view of the age-rotation-activity relation depended until recently on observations of handfuls of stars in the field and in a small number of well-studied clusters, with the Hyades being a particularly key cluster \\citep[e.g.,][]{radick1987, jones1996, stauffer1997, terndrup2000}. Largely because of the advent of time-domain surveys, with their emphasis on wide-field, automated, high-cadence observing, it is now possible to monitor stellar rotation on an entirely new scale \\citep[e.g.,][]{irwin2007, meibom2009, hartman2010}. The Palomar Transient Factory \\citep[PTF;][]{nick2009, rau2009} provides deep, multi-epoch photometry over a wide field-of-view, and our Columbia/Cornell/Caltech PTF (CCCP) survey, one of PTF's Key Projects, is leveraging this capability to measure rotation periods in open clusters of different ages.  Our first CCCP target, Praesepe,\\footnote{Also known as the Beehive Cluster and M44.} 08 40 24 $+$19 41, is a nearby \\citep[$\\sim$180 pc;][]{vanleeuwen2009}, rich \\citep[$\\sim$1200 stars;][]{adam2007}, and intermediate-age \\citep[$\\sim$600 Myr;][]{delorme2011} cluster that shares many characteristics with the Hyades. Until recently, only five rotation periods --- for mid- to late-M dwarfs --- had been measured for Praesepe members \\citep{scholz2007}. These periods were often combined with (less sparse) data for high-mass Hyads in order to infer the mass-rotation relation for 600-Myr-old stars \\citep[e.g.,][]{irwin2009}. This was particularly unsatisfying as the Hyades and Praesepe were the two oldest clusters with measured rotation periods, and were therefore essential in studying the evolution of the age-rotation relation from ages of a few 100 Myr to the age of the Sun.  Fortunately, the situation has improved significantly in the past year. \\citet{delorme2011} surveyed the Hyades and Praesepe as part of the SuperWASP exoplanet-search program. SuperWASP's sensitivity, tuned to discover exoplanets transiting nearby bright stars, enabled the measurement of rotation periods for 52 late-F to late-K/early-M stars in Praesepe. Meanwhile, \\cite{scholz2011} added 49 rotation periods (of which 24 are considered very robust) to the \\citet{scholz2007} sample, with the bulk of this new sample being of spectral type M3-M5.\\footnote{These spectral types are from the \\citet{adam2007} catalog; a few \\citet{scholz2007} and \\citet{delorme2011} stars lack spectral types in this catalog because they are too faint or too bright.} (Thanks to {\\it Kepler}, rotation periods have now also been measured in a 1-Gyr-old cluster, NGC 6811; \\citet{meibom2011b}.)  We report stellar rotation periods for 40 late-K/early-M Praesepe members derived from our first season of PTF observations. Our campaign produced $\\sim$200 distinct observations of four overlapping fields designed to include a large number of Praesepe members identified by \\citet{adam2007}. In Section~\\ref{observations} we describe our data and in Section~\\ref{periods} our period-finding algorithm. We also compare our periods with those derived by \\citet{scholz2007}, \\citet{delorme2011}, and \\citet{scholz2011} for the stars for which they also measured $P_{rot}$, and flag potential binary systems among our rotators. In Section~\\ref{results}, we combine our Praesepe data with that of \\citet{scholz2007}, \\cite{delorme2011}, and \\citet{scholz2011}, and compare color-period relations and mass-period distributions derived from these data to those derived from the Hyades \\citep[using data from][]{delorme2011}, the 150-Myr-old clusters M35 and NGC 2516 (\\citet{meibom2009} and \\citet{irwin2007}, respectively), and kinematically selected young and old field star populations \\citep{kiraga2007}, as well as to gyrochrones derived from the models of \\citet{barnes2010}. We conclude in Section~\\ref{concl}. The Appendix lists interesting variable field stars identified in our Praesepe observations.  In a forthcoming companion paper we use the results of our spectroscopic campaign with the 2.4-m Hiltner telescope at MDM Observatory and the WIYN 3.5-m telescope at NOAO, both on Kitt Peak, AZ, to examine the relationship between rotation and activity in Praesepe.  \\begin{figure}[h] \\centerline{\\includegraphics[angle=90,width=\\columnwidth]{f0.eps}} \\caption{CMD of Praesepe members with 8~$< r_{\\rm SDSS} <$~21 mag. Photometric errors are plotted but are typically smaller than the symbols. Overplotted in red are members with PTF detections; the solid line is the cluster main sequence. Stars with $r \\sim 14 - 16$ and $(r-K)\\ \\gapprox\\ 4$ are ones for which the SDSS photometry is saturated, and for which accurate photometry could not be calculated from the wings of the point spread function.} \\label{detections} \\end{figure} ",
        "conclusions": "Using PTF observations of four overlapping fields, we have measured rotation periods for 40 late-K to mid-M stars belonging to the nearby, rich, $\\sim$600 Myr open cluster Praesepe. Our data occupy a unique area in mass-period space: we measure $P_{rot}$ for stars later than those identified as rotators by \\citet{delorme2011} and generally earlier than those identified by \\citet{scholz2007} and \\citet{scholz2011}, filling the gap between solar-type and mid-to-late M dwarf rotators in Praesepe. Furthermore, the cadence and time-span of our observations gives us sensitivity to the fast rotators identified by these groups and to the slow rotators identified at the high-mass end by \\citet{delorme2011}.  Our measurements indicate that Praesepe's mass-period relation undergoes a transition from a well-defined singular relation to a more scattered distribution of both fast and slow-rotators at masses $\\sim$0.6 $\\Msun$, corresponding roughly to a spectral type of M1. The location of this transition is broadly consistent with expectations based on observations of younger clusters and the assumption that stellar-spin down is the dominant mechanism influencing angular momentum evolution at $\\sim$600 Myr.  A comparison to the data recently published by \\citet{delorme2011} for the Hyades, widely assumed to be coeval to Praesepe, suggests that this transition occurs at different characteristic masses in the two clusters, providing further evidence that Praesepe is the younger of the two clusters. Furthermore, by using the \\citet{barneskim2010} and \\citet{barnes2010} formalisms to evolve the Praesepe $P_{rot}$ in time and comparing the predicted $P_{rot}$ with measured $P_{rot}$ in M35 and NGC 2516 ($\\sim$150 Myr) and for young and old field star populations (1.5 and 8.5 Gyr), we find that stellar spin-down may progress more slowly than described by these relations.  \\begin{deluxetable*}{lccccl}[t!] \\tablewidth{0pt} \\tablecaption{Other Interesting Variable Stars} \\tablehead{ \\colhead{}       & \\colhead {Ave.\\ PTF}    & \\colhead{\\# of}    & \\colhead{$P_{rot}$}   & \\colhead{}   & \\colhead{}\\\\ \\colhead{SDSS J} & \\colhead{$R$ (mag)} & \\colhead{Obs.} & \\colhead{(d)} & \\colhead{Power}  & \\colhead{Type} } \\startdata 083113.96$+$194951.5 & $15.99\\pm0.01$ & 327 & 0.26 & 136.59 & Eclipsing binary? \\\\ 083238.89$+$210424.6\\tablenotemark{a} & $17.44\\pm0.03$ & 173 & 0.55 & 57.20 & RR Lyrae \\\\ 083426.37$+$202040.9 & $15.53\\pm0.01$ & 348 & 0.13 & 60.69 & W Uma? \\\\ 083525.00$+$194659.7 & $15.96\\pm0.01$ & 170 & 0.80 & 63.81 & RR Lyrae \\\\ 083706.80$+$185556.2 & $16.14\\pm0.01$ & 348 & 0.30 & 145.78 & W Uma? \\\\ 083816.78$+$182724.6\\tablenotemark{a} & $17.36\\pm0.03$ & 348 & 0.29 & 108.32 & RR Lyrae \\\\ 084041.70$+$201612.1 & $16.33\\pm0.02$ & 377 & 0.13 & 123.29 & W Uma? \\\\ 084214.42$+$195819.7 & $16.33\\pm0.01$ & 531 & 0.25 & 197.15 & Eclipsing binary? \\\\ 084803.57$+$191340.5 & $16.21\\pm0.01$ & 180 & 0.49 & 55.37 & Known RR Lyrae\\tablenotemark{b}\\\\ 084920.49$+$182617.3 & $15.85\\pm0.01$ & 181 & 0.53 & 63.28 & Known RR Lyrae\\tablenotemark{a,b}\\\\ 084950.50$+$193232.0 & $14.70\\pm0.00$ & 337 & 0.18 & 137.16 & W Uma? \\\\ 085112.62$+$184344.3 & $15.06\\pm0.00$ & 181 & 0.19 & 84.19 & W Uma? \\enddata \\tablenotetext{a}{Spectrum is available from SDSS.} \\tablenotetext{b}{Classified as RR Lyrae in SIMBAD.} \\label{ap_table} \\end{deluxetable*}  The fixed age mass-period relation is but one projection of the underlying stellar age-rotation-activity relationship. Previous studies of stellar activity, in clusters and the field, have derived relationships between a star's age and observational tracers of its coronal or chromospheric activity \\citep[e.g.,][]{skumanich72, radick1987, soderblom2001}. Recent studies of chromospheric activity in low-mass field stars have inferred activity lifetimes as a function of spectral type by modeling the vertical gradient in H$\\alpha$ emission strengths as a consequence of dynamical heating in the Galactic disk \\citep[e.g.,][]{andy08}. These studies predict that stars with spectral types of M2 or later have activity lifetimes $>$1 Gyr. The activity lifetimes of M0-M1 stars are somewhat less well known, as few active early M stars are observed in the field, but appear to be $\\lesssim$600 Myr.  These relations would thus predict that the boundary between H$\\alpha$ active and inactive Praesepe members should occur in the M0/M1 spectral range. The agreement between the implied mass of the active/inactive boundary in Praesepe, near 0.6 $\\Msun$, and the similar characteristic mass for the transition from a singular mass-period relation to a more scattered distribution of rapid and slowly rotators, strengthens the case for an underlying rotation-activity relation in this cluster. In a forthcoming paper, we use the results of our spectroscopic campaign with the 2.4-m Hiltner telescope at MDM Observatory and the WIYN 3.5-m telescope at NOAO, both on Kitt Peak, AZ, to examine this relationship between rotation and activity in Praesepe."
    },
    "1107/1107.4513_arXiv.txt": {
        "abstract": "In this paper we review recent progress in our understanding of the chemical evolution of protoplanetary disks. Current observational constraints and theoretical modeling on the chemical composition of gas and dust in these systems are presented. Strong variations of temperature, density, high-energy radiation intensities in these disks, both radially and vertically, result in a peculiar disk chemical structure, where a variety of processes are active. In hot, dilute and heavily irradiated atmosphere only the most photostable simple radicals and atoms and atomic ions exist, formed by gas-phase processes. Beneath the atmosphere a partly UV-shielded, warm molecular layer is located, where high-energy radiation drives rich ion-molecule and radical-radical chemistry, both in the gas phase and on dust surfaces. In a cold, dense, dark disk midplane many molecules are frozen out, forming thick icy mantles where surface chemistry is active and where complex polyatomic (organic) species are synthesized. Dynamical processes affect disk chemical composition by enriching it in abundances of complex species produced via slow surface processes, which will become detectable with ALMA. ",
        "introduction": "Protoplanetary disks (PPDs) are ubiquitous around young stars \\citep[e.g.,][]{LP80,Lissauer_87a}. Their chemical composition and physical properties regulate the efficiency and timescale of planet formation. Molecules and dust serve as heating and cooling agents of the gas, while dust grains also dominate the disk opacities. Molecular lines observed at infrared and (sub-)millimeter wavelengths are useful probes of physical conditions (temperature, density, kinematics, turbulence) in PPDs.  Over the past decade significant progress has been achieved in our understanding of disk chemical composition, both theoretically and observationally. Since molecular hydrogen is not observable when it is cold, we have to rely on other trace species to infer disk physics and chemistry (Table~\\ref{obs_methods}). Multi-molecule, multi-transition interferometric observations, coupled to line radiative transfer and chemical modeling, allowed to constrain disk sizes, kinematics, distribution of temperature, surface density, and molecular column densities (see reviews by \\citet{Bergin_ea07} and \\citet{DGH07}).  Observations of thermal dust emission are employed to constrain dust properties, infer disk masses, whereas IR spectroscopy is used to study mineralogical composition of dust in PPDs. Recently, with space-borne ({\\it Spitzer}) and ground-based (Keck, VLT, Subaru) infrared telescopes, molecules have been detected in very inner zones of planet-forming systems, at $r \\lesssim 1-10$~AU.  The conditions of planets formation in the early Solar system have been revealed by a detailed analysis of chemical and mineralogical composition of meteoritic samples and cometary dust particles \\citep[e.g.,][]{Bradley_05}. The recent {\\it Stardust} and {\\it Genesis} space missions have returned first samples of pristine materials, likely of cometary origin, showing a complex structure of high-temperature crystalline silicates embedded in low-temperature condensates \\citep{Brownlee_ea04,Flynn_ea06,Brownlee_ea08}. The presence of crystalline silicates in outer regions of protoplanetary disks has also been revealed \\citep[e.g.,][]{vanBoekel_ea04,Juhasz_ea10a}. An isotopic analysis of refractory condensates in unaltered chondritic meteorites shows strong evidence that the inner part of the Solar Nebula has been almost completely mixed during the first several Myr of evolution \\citep[e.g.,][]{Boss2004,Ciesla_09}.  These intriguing findings are partly understood in modern astrochemical models of protoplanetary disks \\citep{ah1999,Mea02,vZea03,TG07,Agundez_ea08,Woods_Willacy08,Visser_ea09,Walsh_ea10,Semenov_Wiebe11a}. The major result of the chemical modeling is that disks have a layered chemical structure due to heavy freeze-out of gas-phase molecules in the cold midplane and their photodissociation in the atmosphere. Observed column densities of CO, HCO$^+$, N$_2$H$^+$, CN, HCN, HNC, CS, etc. are qualitatively agree with the chemical models.  In this paper we briefly review the major observational findings and modeling predictions for the chemical composition and evolution of protoplanetary disks. ",
        "conclusions": ""
    },
    "1107/1107.4808_arXiv.txt": {
        "abstract": "We present a detailed analysis of the partial coverage of the \\qso\\,($z_{\\rm em}=2.57$) broad line region  by an intervening H$_2$-bearing cloud at $z_{\\rm abs}=2.3377$. Using curve of growth analysis and line profile fitting, we demonstrate that the H$_2$-bearing component of the cloud covers the QSO intrinsic continuum source completely but only part of the {Broad Line Region (BLR)}. We find that only 48$\\pm$6~\\% of the C~{\\sc iv} BLR emission is covered by the C~{\\sc i} absorbing gas. We observe residual light ($\\sim6$\\%) as well in the bottom of the O~{\\sc i}\\,$\\lambda$1302 absorption from the cloud, redshifted on top of the QSO Lyman-$\\alpha$ emission line. Therefore the extent of the neutral phase of the absorbing cloud is not large enough to cover all of the background source. The most likely explanation for this partial coverage is the small size of the intervening cloud, which is comparable to the BLR size. We estimate the number densities in the cloud: $n_{\\rm H_2}$~$\\sim$~110~cm$^{-3}$ for the \\mbox{H$_2$-bearing} core and $n_{\\rm H}$~$\\sim$~30~cm$^{-3}$ for the neutral envelope. Given the column densities, $N$(H$_2$)~=~3.71$\\pm$0.97$\\times$10$^{19}$~cm$^{-2}$ and $N$(H~{\\sc i})~=~7.94$\\pm$1.6$\\times$10$^{20}$~cm$^{-2}$, we derive the linear size of the {H$_2$-bearing} core and the neutral envelope along the line of sight to be $l_{\\rm H_2}$~$\\sim$~$0.15^{+0.05}_{-0.05}$~pc and $l_{\\rm HI}$~$\\sim$~$8.2^{+6.5}_{-4.1}$~pc, respectively. We estimate the size of the C~{\\sc iv} BLR by two ways (i) extrapolating size$-$luminosity relations derived from reverberation observations and (ii) assuming that the H$_2$-bearing core and the BLR are spherical in shape and the results are $\\sim$0.26 and $\\sim$0.18~pc, respectively. The large size we derive for the extent of the neutral phase of the absorbing cloud together with a covering factor of $\\sim$0.94 of the Lyman-$\\alpha$ emission means that the Lyman-$\\alpha$ BLR is probably fully covered but that the Lyman-$\\alpha$ emission extends well beyond the limits of the BLR. ",
        "introduction": "\\label{introduction} \\noindent The broad emission lines in the spectra of active galactic nuclei respond to variations in the luminosity of the central continuum source with a delay due to light-travel time effects within the emission-line region. It is therefore possible through the process of 'reverberation mapping' to determine the geometry and kinematics of the emission-line region by careful monitoring of the continuum variations and the resulting emission-line response (Blandford \\& McKee 1982; Peterson 1993; Netzer \\& Peterson 1997). In particular the size of the broad line region (BLR) can be inferred {from the time delay measurement}. Recent investigations of low-redshift AGNs show a tight relation between this size and the luminosity of the AGN, $R$~=~$A\\times$($L$/10$^{43}$)$^{B}$, where $R$ is the radius of the BLR, $A$ is a typical distance in light-days and $L$ is the H$\\beta$ luminosity in erg/s. The index is found to have a value close to $B$~$\\sim$~0.6-0.7 when the typical distance $A$ is in the range 20-80 light-days (Wu et al. 2004; Kaspi et al. 2005). Extending this relation to high luminosities yields a typical radius of the order of 1~pc for the BLR of bright high-$z$ quasars. The size of the BLR has also been shown to be correlated with the luminosity in the continuum (Bentz et al. 2009).  The anti-correlation found between the radius of the region over which an emission line is emitted and the velocity width of the broad emission line in the same object supports the idea that the BLR gas is virialized and its velocity field is dominated by the gravity of the central black-hole (Peterson \\& Wandel 1999). If this is the case, then the BLR size and the emission line width give an estimate of the mass of the central object (Peterson \\& Wandel 1999; Warner, Hamann \\& Dietrich 2003; Wang et al. 2009). The broad line region is stratified and the BLR  reverberation mapping size for C~{\\sc iv} is about half that for H$\\beta$. This is consistent with the above assumption as more highly ionized species are expected to be found primarily closer to the central source of ionization radiation.  The spatial extent of the BLR is revealed by the partial coverage of some absorbing clouds, usually associated with the AGN, located in front of the quasar and producing absorption lines that are saturated but do not go to the zero flux level. Usually, the continuum source is covered completely but the emission line region can be covered only partially (e.g. Petitjean, Rauch \\& Carswell 1994, Hamann 1997, {Srianand \\& Shankaranarayanan 1999}). In Wampler, Chugai \\& Petitjean (1995), four Fe~{\\sc ii} clouds are seen at different velocities with the similar covering factor, $f\\sim 0.5$. In Srianand et al. (2002), line locking and covering factors are shown to be intimately related and are used to constrain the geometry of the BLR. Covering factor is one of the characteristics together with variability and high metallicity that are used to distinguish intrinsic from intervening absorption systems. Indeed, partial coverage of intervening systems has rarely been reported. It has been the case in the early Keck spectrum of APM08279+5255 (Ellison et al. 1999; Petitjean et al. 2000a) which is a lensed quasar whose images are separated by only 0.35~arcsec so that the Keck spectrum encompasses all the images (Ledoux et al. 1998). It is the case that the intervening Mg~{\\sc ii} systems are not covering all the lensed images (Lewis et al. 2002; Ellison et al. 2004) and, because of this, typical dimensions of the intervening clouds are derived to be of the order of $\\sim$1~kpc.  Partial coverage of a BLR by an intervening absorber had never been reported before Ivanchik et al. (2010). These authors note that the C~{\\sc i} lines associated with the $z_{\\rm abs}$~=~2.3377 DLA system towards \\qso\\,\\, probably do not cover the C~{\\sc iv} BLR completely so that some flux stays unabsorbed at the bottom of saturated lines. In the present paper, we analyse in details this unique effect and test different interpretations. We present the observations in Section~2. Partial coverage is ascertained in Section~3. Physical conditions of the gas in the DLA are derived in Section~4 in order to infer {its} extent. Results are discussed in Section~5 before conclusions are drawn in Section~6. ",
        "conclusions": "\\noindent \\label{conclusion}  The analysis of H$_2$, C\\,{\\sc i}, O~{\\sc i} and Si~{\\sc ii} absorption lines from the molecular DLA system at $z_{\\rm abs}$~=~2.3377 toward \\qso\\, has shown that the intervening absorbing cloud is not covering the background source totally. We used a curve-of-growth and a profile fitting analysis to estimate the partial covering factor. The different methods yield covering factors of the {H$_2$-bearing} core (as traced by H$_2$ and C~{\\sc i}) to be $\\approx$~48, 66 and 75~\\% for the C~{\\sc iv}, C~{\\sc iii} and Lyman-$\\beta$-O~{\\sc vi} emission lines {whilst} the QSO intrinsic continuum is covered completely. The O~{\\sc i}\\,$\\lambda$1302 and Si~{\\sc ii}\\,$\\lambda$1304 absorptions cover only $\\sim$94\\% of the Lyman-$\\alpha$ emission.  According to the generally accepted model of AGNs, broad emission lines are emitted by warm and highly ionized gas located in the BLR with transverse dimension of the order of $\\sim$1~pc. The quasar continuum is produced by the inner part of the accretion disk, which linear size is several orders of magnitude less than the BLR size. Thus the most probable explanation of the observed partial coverage is the comparable angular sizes of the BLR and the compact {H$_2$-bearing} absorption cloud. The fact that the continuum is completely absorbed makes other explanations such as the presence of a binary quasar or gravitational lensing less plausible.  We derived the linear extent of the H$_2$-bearing cloud and neutral envelope to be $l_{\\rm H_2}$~$\\sim$~$0.15^{+0.05}_{-0.05}$~pc and $l_{\\rm HI}$~$\\sim$~$8.2^{+6.5}_{-4.1}$~pc, respectively. Assuming that the {H$_2$-bearing} component is spherical in shape, we estimate the size of the C~{\\sc iv} BLR to be $\\sim$0.2~pc. The large size we derive for the neutral cloud extent together with the covering factor of $\\sim$94\\% of the Lyman-$\\alpha$ emission means that the Lyman-$\\alpha$ BLR is probably fully covered but that the Lyman-$\\alpha$ emission extends well beyond the extent of the C IV emitting BLR. In turn this means that the extended Lyman-$\\alpha$ emission corresponds to about 6~\\% of the flux at the peak of the emission (over $\\sim$500~km/s). Assuming the C~{\\sc iv} BLR is virialized, we derive the mass of the central $M_{\\rm BH}$~=~6.8$^{+4.1}_{-4.5}$$\\times$10$^{8}$~M$_{\\odot}$  Partial coverage of the background source by intervening clouds has been observed an used to derive the radius of the clouds when the background source is the combination of several gravitational images (Petitjean et al., 2000a; Ellison et al., 2004). This is the first time that partial coverage of a QSO BLR by an intervening cloud is reported.  This kind of situation tests the kinematics of the BLR. Indeed we see that the covering factor depends of the position of the absorption on top of the emission line. This can be due to the presence of inflows or outflows in BLR (Denney et al. 2009) which could be covered differently or the BLR could have a disk-like structure implying the projected velocity is correlated with the spatial position. Obviously it will be difficult to {disentangle} the different kinematics. However, it may be interesting to dramatically increase the quality of spectra in which an intervening compact absorption system is observed. This may reveal that partial covering factor is not unusual and could be used to constrain the BLR structure.   \\vspace{2mm}{\\footnotesize {\\rm Acknowledgments.} This work was supported in part by a bilateral program of the Direction des Relations Internationales of CNRS in France, by the Russian Foundation for Basic Research (grant 11-02-01018a), and by a State Program ``Leading Scientific Schools of Russian Federation'' (grant NSh-3769.2010.2), by Ministry of Education and Science of Russian Federation (contract \\# 11.G34.31.0001 with SPbSPU and leading scientist G.G. Pavlov). PPJ and RS acknowledge support from the Indo-French Centre for the Promotion of Advanced Research under the programme No.4304--2. SB thanks Dynasty Foundation and A.M. Krassilchtchikov for help with cluster computations.}"
    },
    "1107/1107.3872_arXiv.txt": {
        "abstract": "Using a vector vortex coronagraph behind the 1.5-m well-corrected subaperture (WCS) at Palomar, we detected a second object very close to $\\epsilon$~Cephei, a $\\delta$~Scuti F0 IV star. The candidate companion, $\\sim 50$ times fainter than  $\\epsilon$~Cephei, if physically associated, is a late-type K or early M star, and lies at an angular separation of 330 mas, or 1.1 $\\lambda/D$ for the WCS, making it the smallest angle detection ever realized with a coronagraph in terms of $\\lambda/D$ units. The projected separation of the putative companion is $\\sim 8.6$ AU, most likely on a highly eccentric orbit. The recently detected near-infrared excess is thus likely not due to hot dust. Moreover, we also show that the previously reported {\\it IRAS} 60 $\\mu$m excess was due to source confusion on the galactic plane. ",
        "introduction": "High contrast imaging at small angles is a very useful tool to access inner regions around stars. While instrumental limits are constantly improved by the development of exoplanet imaging and characterization techniques, such capabilities can be used in stellar astrophysics to discover new systems or put constraint on known ones. Moreover, this small-angle capability allows a significant reduction in the size of potential space telescopes aimed at detecting and characterizing exoplanets, which is a cost-effective way of building exoplanet missions.  We have used a 1.5-m subaperture of the 5.1-m Hale telescope at Palomar with a small inner working angle (IWA) phase-mask coronagraph to detect a new candidate stellar companion to the $\\delta$~Scuti star $\\epsilon$~Cephei. The result is interesting for two reasons. First, $\\epsilon$~Cephei was claimed to have an infrared excess measured by {\\it IRAS}, and recently detected by near-infrared ground-based interferometry. Such infrared excesses are generally attributed to dusty debris disks \\citep{Absil2006}. We show here with an image that the origin of the near-infrared excess is a second object that is likely a late-type stellar companion. Moreover, we demonstrate that the previously reported 60 $\\mu$m {\\it IRAS} excess is attributable to an unrelated field object. Second, the potential companion was imaged at 1.1 diffraction beamwidths ($\\lambda/D$, with the working wavelength $\\lambda\\simeq 2.16$ $\\mu$m, and the telescope diameter $D=1.5$ m) from the star. Such a small angle detection was only possible because of our phase mask vector vortex coronagraph (VVC), demonstrating that such coronagraphs can effectively provide very small IWA. ",
        "conclusions": "This letter reports the detection of a candidate dim companion to $\\epsilon$~Cephei, a $\\delta$~Scuti star. This is the first image of a candidate companion, seen at 330 $\\pm 50$ mas separation, only 1.1 resolution elements from the central host star ($\\sim 8.6$ AU). If physically associated, the companion has a mass of $\\sim 0.5$ $M_{\\rm Sun}$, and a temperature of 3650 K, likely placing it at the transition between late-type K stars and early M stars. The interferometrically detected near-infrared excess is likely due to the companion, and not hot dust, while the previously reported {\\it IRAS} 60 $\\mu$m excess was due to source confusion. Finally, the most plausible reason why this candidate companion remained hidden so far is that {\\it Hipparcos} failed to detect it because of its long period, highly eccentric orbit, while the stellar pulsations of the $\\delta$~Scuti primary precluded an RV detection.    This detection was made possible by state-of-the-art high contrast imaging techniques (wavefront control and vortex coronagraphy) envisioned for upcoming next-generation ground-based extreme AO systems and future space-based coronagraphs. We have demonstrated a $5\\sigma$ detection capability of the order of $\\sim 2 \\times 10^{-3}$ at $\\sim 1\\lambda/D$, which is sufficient to directly image secondary stars and brown dwarfs to an order of magnitude fainter than speckle interferometry. At $2\\lambda/D$, we routinely reach $10^{-4}$ contrast levels, and with careful calibration $\\sim 10^{-5}$ \\citep{Mawet2010b,Serabyn2010}, two to three orders of magnitude better than current interferometric techniques. Extreme AO coupled to vector vortex coronagraphy has thus great potential for the characterization of spectroscopic binary and brown dwarf orbits, especially with a larger telescope. It can also dramatically improve imaging capabilities even with rather small telescopes. Such techniques thus open the door to the direct imaging of spectroscopic binary companions, as well as brown dwarfs and exoplanets."
    },
    "1107/1107.3261_arXiv.txt": {
        "abstract": "We report on the $\\lambda = 5-36 \\mu$m {\\it Spitzer} Infrared Spectrograph spectra of 79 young stellar objects in the very young nearby cluster NGC 1333. NGC 1333's youth enables the study of early protoplanetary disk properties, such as the degree of settling as well as the formation of gaps and clearings. We construct spectral energy distributions (SEDs) using our IRS data as well as published photometry and classify our sample into SED classes. Using ``extinction-free\" spectral indices, we determine whether the disk, envelope, or photosphere dominates the spectrum. We analyze the dereddened spectra of objects which show disk dominated emission using spectral indices and properties of silicate features in order to study the vertical and radial structure of protoplanetary disks in NGC 1333. At least nine objects in our sample of NGC 1333 show signs of large (several AU) radial gaps or clearings in their inner disk. Disks with radial gaps in NGC 1333 show more-nearly pristine silicate dust than their radially continuous counterparts. We compare properties of disks in NGC 1333 to those in three other well studied regions, Taurus-Auriga, Ophiuchus and Chamaeleon I, and find no difference in their degree of sedimentation and dust processing. ",
        "introduction": "\\label{intro} NGC 1333 was first identified as an optical reflection nebula at the western end of the Perseus molecular cloud. The moniker NGC 1333 is also used to identify the molecular cloud and young stellar cluster associated with the optical reflection nebula along with the dark cloud L1450 \\citep{herbig72, herbig74, strom76, walawender08}. NGC 1333 is one of the best studied very young clusters of active low- to intermediate-mass star formation \\citep{aspin94, lada96, wilking04}. Imaging surveys in the near- and mid-infrared have revealed NGC 1333 contains both a northern and southern young stellar cluster \\citep{aspin94, lada96, gut08}. Previous studies have also shown NGC 1333 is heavily extinguished and contains a substantial fraction of deeply embedded young stellar objects (YSOs) as well as many prominent outflows \\citep{lada96,knee00,winston09}. The {\\it Spitzer Space Telescope} \\citep{werner04} infrared view of NGC 1333 in \\citet{gut08} is streaked with emission at 4.5$\\mu$m, which is dominated by H$_2$ lines  arising in shocks from YSO outflows colliding with ambient cloud material. These energetic outflows originate from actively accreting protostars with infalling envelopes (Class 0, I). As they are surrounded by dusty envelopes, these protostars as well as more evolved pre-main-sequence stars with circumstellar disks (Class II) emit in the mid-infrared and are well detected by {\\it Spitzer}. Pre-main-sequence stars with little to no dusty disk or envelope (Class III) are not as easily detected in the mid-infrared and have been found in NGC 1333 through X-ray surveys \\citep{getman02,winston09,winston10}.  Many previous studies have addressed the age of NGC 1333. In a near-infrared (NIR) spectrographic and X-ray study of NGC 1333, \\citet{winston09} estimated most objects in NGC 1333 to be $<$3 Myr, and found an age spread from 1 to 10 Myr based on isochrones by \\citet{baraffe98}. An optical and NIR study of NGC 1333 by \\citet{aspin03} found similar results using isochrones by \\citet{dantona97} namely an age spread from 7$\\times$10$^4$ yr to 10 Myr with many sources around 3 Myr. They interpreted this spread in ages to be evidence of multiple epochs of star formation potentially spurred by turbulence generated by the many prominent outflows in NGC 1333. Others have also suggested turbulence driven star formation could have taken place in NGC 1333 \\citep[e.g.][]{quillen05}. Studies of the low mass (0.25 - 0.02 M$_{\\odot}$) population of NGC 1333 estimate a median age of $\\sim$0.3 Myr for that population using the low mass stellar evolution tracks by \\citet{dantona97} \\citep{greissl07, wilking04}. However, for objects less than 1 Myr old age estimates based on isochrones are more uncertain than for older objects as the location of the birth line (t = 0) is ambiguous \\citep{dantona97, siess01}. For this reason, class fractions are a common tool used to get an idea of the relative youth of a cluster. \\citet{gut08} show NGC 1333 is very young due to its large disk fraction (83\\%$\\pm$11\\%) which they define as the fraction of total objects which show a large infrared excess at certain wavelengths indicative of emission from a circumstellar disk. \\citet{lada96} estimate the age of NGC 1333 to be very young based on the similarity of its $K$-band luminosity function to that of the Trapezium ($\\la$ 1 Myr), as well as its large fraction of infrared sources (61\\%) which suggest an age of $\\le$ 1-2 Myr. The protostar and Class II populations also have the same peak nearest neighbor distance (0.045 pc), meaning stars do not appear to have moved far from where they formed, another indication of youth \\citep{gut08}. By these measures, NGC 1333 is one of the youngest clusters observed comprehensively by {\\it Spitzer}. Based on disk fraction, NGC 1333 is younger than three well-studied star forming regions also observed by {\\it Spitzer}, namely Taurus ($\\sim$ 1-2 Myr; \\citet{furlan06}), Ophiuchus ($\\sim$ 0.8 Myr; \\citet{mcclure10}), and Chamaeleon I ($\\sim$ 2 Myr; \\citet{manoj11}). We adopt an age of $\\la$ 1 Myr for this study.  The age of a cluster, while difficult to determine, is necessary to understand the timescales of star formation and disk evolution processes. Observations show disk dissipation is well underway in 3-5 Myr, as on this timescale infrared excess around young stellar objects drops by half \\citep{haisch01,her...dez08} and accretion rate falls below 10$^{-10}$ M$_{\\odot}$ yr$^{-1}$ for low mass stars \\citep[e.g.][]{muzerolle03}. Therefore the study of regions significantly younger than this allow for investigation of early disk evolution and dissipation processes and can further constrain the timescales of these phenomena. Circumstellar dust and gas predominately emit in the infrared making it the ideal wavelength range to investigate processes in disk evolution. Models predict \\citep{weidenschilling97,harker02, gail04,dullemond05,dalessio06} and observations confirm \\citep{furlan05,kessler06,sicilia07, sicilia08, sicilia09,bouwman08, sargent09, olofsson09, olofsson10} that grain growth, dust processing and dust settling in disks proceed on timescales $<$1 Myr.  Giant planets are predicted to form within a few million years through collisional growth of grains into multi-M$_\\oplus$ cores which may then accrete gas from the disk \\citep{pollack96,weidenschilling08, dodson08}. They may also form through gravitational instability of gas and dust in the disk in $\\la$10$^5$yr \\citep{gold73,boss01}. The gravitational interaction between a disk and a giant planet would clear out material in its orbit thereby opening a gap or forming an inner clearing \\citep[e.g.][]{borderies89}. Disks with radial gaps and holes have been detected mostly via deficits in their spectral energy distributions (SEDs) compared to the median for their population \\citep{strom89,rice03b, dalessio05, brown07, calvet05,espaillat07,espaillat08, kim09, merin10}, but in several cases have been confirmed by $mm$-wave interferometric imaging \\citep[e.g.][]{pietu06, hughes09}. The leading explanation for such gaps and clearings in disks is the gravitational and tidal effect of companions with the central star. In the paradigm ``central clearing\" disk around CoKu Tau/4 the companion is a star \\citep{ireland08}, but in most others it is likely to be a gas-giant planet \\citep{artymowicz94,calvet05,kim09,pott10}. However, several disk-dissipation mechanisms which clear disks from the inside out have been invoked to explain central clearings, among them photoevaporation \\citep{clarke01, alexander06} and MRI draining \\citep{BH91, chiang07}. These methods generally fail to account for gaps separating inner and outer disks, which outnumber the central clearings \\citep{furlan09,kim09,mcclure10,manoj11}. Disks with holes or gaps are historically called transitional disks (TDs) as they appear to be in transition between a star with a full disk of dust and gas around it (Class II object) and a star where little to no emission is seen from a disk(Class III object) \\citep{strom89, skrutskie90,gauvin92}. Disks with radial gaps are a type of transitional disk known as `pre-transitional disks' \\citep[PTDs;][]{espaillat07, espaillat08,espaillat09, espaillat10}. The interaction of one or more planets with the disk is key to understanding the timescales and processes of planet formation as well as disk dissipation.  Here we present the spectra of 81 objects in NGC 1333. In \\S~\\ref{obs} we discuss the observations, selection, completeness, and data reduction for our sample. We correct our YSO sample for interstellar extinction in \\S~\\ref{EC} and using their observed spectral indices we classify our sample into SED classes and evolutionary states in \\S~\\ref{sam}. In \\S~\\ref{Anal} we characterize the settling in different regions of our disks as well as identify objects with interrupted radial disk structure. We also compare the distribution of various spectral properties and the median spectrum of disks in NGC 1333 to those in Taurus, Ophiuchus and Chamaeleon I. Further discussion of the classification of TDs and PTDs in NGC 1333 is found in \\S~\\ref{disc}. Finally, in \\S~\\ref{conclusion} we summarize our results and state our conclusions. ",
        "conclusions": "\\label{conclusion} We presented the IRS spectra of 79 YSOs in NGC 1333. We supplemented the IRS spectra with published photometry and constructed SEDs. We used the SEDs, before extinction correction, to classify our sample into SED classes as well as evolutionary states, using ``extinction-free\" indices. Based on this classification, our sample in NGC 1333 consists of 21 objects with envelope dominated emission, 57 objects with disk dominated emission and 1 object with photosphere dominated emission. Further, we analyzed various spectral properties of 44 objects which show disk dominated emission and had no evidence of an envelope in their observed SEDs. We compared the shape of the continuum emission as well as silicate emission features to models of radially continuous irradiated accretion disks in an effort to understand the vertical and radial structure of these disks in NGC 1333. We used two different methods in order to identify disks with altered radial structure. With these analyses, as well as comparison with a representative median spectrum, we classified TD and PTDs in our sample. We also compared the properties of disks in NGC 1333 to disks in the Taurus, Ophiuchus, and Chamaeleon I star forming regions in order to investigate potential differences with age and population. Our main results are summarized below:  \\begin{itemize} \\item We find no difference between the distribution of dust settling continuum indices, $n_{6-13}$ and $n_{13-31}$ in NGC 1333, Taurus, L1688 in Ophiuchus, and Chamaeleon I. However, through comparison of class fractions, we find NGC 1333 is likely to be even younger than the youngest of those three regions, Ophiuchus. This similarity between star forming regions of significantly different ages implies advanced dust settling ($\\epsilon$ $\\sim$ 0.001) can happen very early in the life of a disk. Alternatively, few disks are well mixed even at the young age of $\\lesssim$ 1 Myr.  \\item Of the disks analyzed in NGC 1333, 20\\% show evidence of radial gaps and clearings and 11\\% are identified as transitional disks. These fractions are significantly larger than those of Taurus, Ophichus and Chamaeleon I. We find these observations of NGC 1333 to be evidence of giant planet formation resulting in disk dissipation at a very young age, $\\lesssim$ 1 Myr.  \\item Most PTDs, which are large outliers in EW$_{10}$, were not large outliers in EW$_{20}$. This suggests the gaps in these disks which produce an enhanced EW$_{10}$ are not large enough to affect EW$_{20}$. The smaller gaps in these PTDs likely results in an $n_{13-31}$ similar to radially continuous disks and makes them difficult to identify in analyses which use a high $n_{13-31}$ as a criterion.  \\item Multiple analyses are needed to identify PTDs from spectra without detailed modeling. The $n_{2-6}$ vs. $n_{13-31}$ was not effective in selecting PTDs in NGC 1333 and in some cases contradicted the classification via EW$_{10}$ vs. $n_{13-31}$ and comparison to a scaled representative median.  \\item Comparison of a representative median of disks in NGC 1333 to Taurus, Chamaeleon and L1688 demonstrates the median spectrum of NGC 1333 had a similar shape to all three regions, though after scaling was still fainter. The prevalence of late type stars in our NGC 1333 sample can account for this flux discrepancy.  \\item PTDs in our sample NGC 1333 have more pristine optically thin dust in the inner parts of their disks based on analysis of the shape of their 10$\\mu$m emission using F$_{11.3}$/F$_{9.8}$. The dust in these disks with radial gaps seems to have experienced less grain growth and crystallization than radially continuous disks.  \\end{itemize}"
    },
    "1107/1107.1863_arXiv.txt": {
        "abstract": "In the context of the geometrical analysis of weakly non Gaussian CMB maps, the 2D differential extrema counts as functions of the excursion set threshold is derived from the full moments expansion of the joint probability distribution of an isotropic random field, its gradient and invariants of the Hessian. Analytic expressions for these counts are given to second order in the non Gaussian correction, while a Monte Carlo method to compute them to arbitrary order is presented. Matching count statistics to these estimators is illustrated on fiducial non Gaussian  ``Planck\" data. ",
        "introduction": "Extrema counts, especially that of the maxima of the field, have long application to cosmology \\cite[e.g.][]{BBKS}, however  theoretical development  have been mostly restricted to the Gaussian fields.  The statistics of extrema counts, as well as of the Euler number, requires the knowledge of the one-point JPDF $P(x,x_i,x_{ij})$ of the field $x$, its first, $x_i$, and second, $x_{ij}$, derivatives \\footnote{\\label{sigmas}We consider the field and its derivatives to be normalized by their corresponding variances $\\sigma_0^2=\\left\\langle x^2 \\right\\rangle, \\sigma_1^2 = \\left\\langle (\\nabla x)^2 \\right\\rangle, \\sigma_2^2 = \\left\\langle (\\Delta x)^2 \\right\\rangle $. This implies that in our dimensionless units $\\langle x^2 \\rangle=\\langle q^2 \\rangle=\\langle J_1^2 \\rangle=\\langle J_2 \\rangle=1$. }. Extrema density is an intrinsically isotropic statistics given by \\citep{Adler,Longuet} \\begin{equation} \\frac{\\partial n_{\\rm ext}}{ \\partial x} = \\int {\\rm d}^6 x_{ij} P(x,x_i=0,x_{ij}) |x_{ij}| \\,. \\label{eq:ext_int} \\end{equation} Under the condition of statistical isotropy of the field, the essential form for the JPDF is therefore given in terms of the rotation invariants --- $x$ itself, the square of the magnitude of the gradient $q^2\\equiv x_1^2+x_2^2$ and the two invariants $J_1\\equiv\\lambda_{1}+\\lambda_{2}$,  $J_2\\equiv(\\lambda_{1}-\\lambda_{2})^2$ of the Hessian matrix $x_{ij}$ (where $\\lambda_i$ are the eigenvalues of the Hessian). Introducing  $\\zeta=(x+\\gamma J_1)/\\sqrt{1-\\gamma^2}$ (where the spectral parameter $\\gamma=- \\langle x J_1 \\rangle$ characterizes the shape of the underlying power spectrum), leads to the following JPDF for the Gaussian 2D field \\begin{equation} G_{\\rm 2D} = \\frac{1}{2 \\pi} \\exp\\left[-\\frac{1}{2} \\zeta^2 - q^2 - \\frac{1}{2} J_1^2 - J_2 \\right] \\,. \\label{eq:2DG} \\end{equation} The invariant form for the extrema counts \\begin{equation} \\frac{\\partial n_{\\rm ext}}{\\partial x} = \\int \\!\\!\\! \\frac{{{\\rm d} J_1} {{\\rm d} J_2}}{8\\pi^2 \\sqrt{1\\!-\\!\\gamma^2}} \\exp\\left[-\\frac{1}{2} \\zeta^2 - \\frac{1}{2} J_1^2 -J_2\\right] \\left|J_1^2-J_2 \\right| \\nonumber \\end{equation} then readily recovers the classical results  \\citep{Adler,Longuet,BBKS} when the limits of integration that define the extrema type are implemented, namely $J_1 \\in [-\\infty,0]$, $J_2 \\in [0, J_1^2] $ for maxima, $J_1 \\in [0,\\infty]$, $J_2 \\in [0, J_1^2] $ for minima and $J_1 \\in [-\\infty,\\infty]$,$J_2 \\in [J_1^2, \\infty]$  for saddle points.  In \\cite{PGP} we have observed that for non-Gaussian  JPDF the invariant approach immediately suggests a Gram-Charlier expansion  in terms of the orthogonal polynomials defined by the kernel $G_{\\rm 2D}$. Since $\\zeta$, $q^2$, $J_1$ and $J_2$ are uncorrelated variables in the Gaussian limit, the resulting expansion is \\begin{widetext} \\begin{equation} P_{\\rm 2D}(\\zeta, q^2, J_1, J_2) =  G_{\\rm 2D} \\left[ 1 + \\sum_{n=3}^\\infty \\sum_{i,j,k,l=0}^{i+2 j+k+2 l=n} \\frac{(-1)^{j+l}}{i!\\;j!\\; k!\\; l!} \\left\\langle \\zeta^i {q^2}^j {J_1}^k {J_2}^l \\right\\rangle_{\\rm GC} H_i\\left(\\zeta\\right) L_j\\left(q^2\\right) H_k\\left(J_1\\right) L_l\\left(J_2\\right) \\right]\\,, \\label{eq:2DP_general} \\end{equation} \\end{widetext} where terms are sorted in the order of the field power $n$ and $\\sum_{i,j,k,l=0}^{i+2 j+k+2 l=n} $ stands for summation over all combinations of non-negative $i,j,k,l$ such that $i+2j+k+2l$ adds to the order of the expansion term $n$.  The Gram-Charlier coefficients, $\\left\\langle \\zeta^i {q^2}^j {J_1}^k {J_2}^l \\right\\rangle_{\\rm GC}\\equiv (-1)^{j+l} j! l! \\left\\langle H_i\\left(\\zeta\\right) L_j\\left(q^2\\right) H_k\\left(J_1\\right) L_l\\left(J_2\\right) \\right\\rangle_{\\rm m}$ that appear in the expansion can be related to the more familiar cumulants of the field and its derivatives (we use $\\langle \\,\\,\\,\\,\\rangle_{\\rm m}$ for statistical moments while reserving $\\langle \\,\\,\\,\\,\\rangle$ for statistical cumulants), actually being identical to them for the first three orders $n=3,4,5$. Lookup tables of the relationship between Gram-Charlier cumulants and statistical  cumulants can be found at {\\tt http://www.iap.fr/users/pichon/Gram/}. As an illustration, one sixth order non trivial cumulant would be $\\left\\langle J_1^3 J_2 \\zeta \\right\\rangle {}_{\\text{CG}}=\\left\\langle J_1^3 J_2 \\zeta \\right\\rangle+\\left\\langle J_1^3\\right\\rangle  \\left\\langle J_2 \\zeta \\right\\rangle +3 \\left\\langle J_1 J_2\\right\\rangle  \\left\\langle J_1^2 \\zeta \\right\\rangle $. It is prudent to stress that the Gram-Charlier series expansion is distinct from the perturbative  expansions. For instance, while the linear Edgeworth or $f_{\\rm NL}$ expansion  match solely to the first order $n=3$ Gram-Charlier coefficients, quadratic terms require knowledge of the Gram-Charlier terms to $n=6$, while the cubic ones to $n=9$.  Integrals over $J_1$ and $J_2$ for extremal points can be carried out analytically even for the general expression (\\ref{eq:2DP_general}). Different types of critical points can be evaluated separately by restraining the integration domain in the $J_1$-$J_2$ plane to ensure the appropriate signs for the eigenvalues.  The effect of the non-Gaussian cubic correction on the total number of the extrema of different types is given by \\begin{eqnarray} n_{\\rm max/min} &=&  \\frac{1}{8 \\sqrt{3} \\pi {R_*}^2} \\pm \\frac{18 \\left\\langle q^2 J_1 \\right\\rangle - 5\\left\\langle J_1^3 \\right\\rangle + 6 \\left\\langle J_1 J_2 \\right\\rangle}{54 \\pi \\sqrt{2\\pi} {R_*}^2} \\,, \\nonumber \\\\ n_{\\rm sad} &=& \\frac{1}{4 \\sqrt{3} \\pi{R_*}^2} \\,, \\end{eqnarray} where we have restored (see note [11]) the dimensional scaling with $R_* = \\sigma_1/\\sigma_2$ , the characteristic separation scale between  extrema. The total number of saddles, as well as of all the extremal points, $n_{\\rm max} + n_{\\rm min} + n_{\\rm sad}$, are preserved in the first order (the latter following for the former, as topological considerations imply $n_{\\rm max}-n_{\\rm sad}+n_{\\rm min}={\\rm const}$), but the symmetry between the minima and the maxima is  broken.  The differential number counts with respect to the excursion threshold $\\nu$ are given by \\begin{widetext} \\begin{eqnarray} \\frac{\\partial n_{\\rm max/min}}{\\partial \\nu} &=& \\frac{1}{\\sqrt{2 \\pi}{R_*}^2} \\exp\\left(-\\frac{\\nu^2}{2}\\right) \\left[1 \\pm \\mathrm{erf}\\left(\\frac{\\gamma \\nu}{\\sqrt{2(1-\\gamma^2)}}\\right) \\right] K_1(\\nu,\\gamma) \\pm \\frac{1}{\\sqrt{2 \\pi (1-\\gamma^2)}{R_*}^2} \\exp\\left(-\\frac{\\nu^2}{2(1-\\gamma^2)}\\right) K_3(\\nu,\\gamma) \\nonumber \\label{eq:difnu1} \\\\ &+& \\frac{\\sqrt{3}}{\\sqrt{2 \\pi (3-2\\gamma^2)}{R_*}^2} \\exp\\left(-\\frac{3 \\nu^2}{6-4 \\gamma^2}\\right) \\left[1 \\pm \\mathrm{erf}\\left(\\frac{\\gamma \\nu}{\\sqrt{2(1-\\gamma^2)(3-2\\gamma^2)}}\\right) \\right] K_2(\\nu,\\gamma) , \\\\ \\frac{\\partial n_{\\rm sad}}{\\partial \\nu} &=& \\frac{2 \\sqrt{3}}{\\sqrt{2 \\pi (3-2\\gamma^2)}{R_*}^2} \\exp\\left(-\\frac{3 \\nu^2}{6-4 \\gamma^2}\\right) K_2(\\nu,\\gamma), \\label{eq:difnu2} \\end{eqnarray} where $K_1, K_2, K_3$ are polynomials with coefficients expressed in terms of the cumulants.  Here we give explicit expressions for the first non-Gaussian order, while the next order can be found at the above mentioned URL.  The term $K_1(\\nu,\\gamma)$ has a special role determining the Euler number $\\chi(\\nu)$ via ${\\partial  \\chi}/{\\partial \\nu} = {\\partial}/{\\partial \\nu} \\left(n_{\\rm max} + n_{\\rm min} - n_{\\rm sad}\\right) = \\sqrt{{2}/{\\pi}} \\exp(-{\\nu^2}/{2}) K_1(\\nu,\\gamma)$. As such, its full expansion has been given in \\cite{PGP}, Eq.~(7), and confirmed to the second order in \\cite{matsu10}. To the leading non-Gaussian order \\begin{equation} K_1 = \\frac{\\gamma^2}{8\\pi} \\left[ H_2(\\nu) + \\left( \\frac{2}{\\gamma} \\left\\langle q^2 J_1 \\right\\rangle + \\frac{1}{\\gamma^2} \\left\\langle x {J_1}^2 \\right\\rangle - \\frac{1}{\\gamma^2} \\left\\langle x J_2 \\right\\rangle \\right) H_1(\\nu) - \\left( \\left\\langle x q^2 \\right\\rangle + \\frac{1}{\\gamma} \\left\\langle x^2 J_1 \\right\\rangle \\right) H_3(\\nu) + \\frac{1}{6} \\left\\langle x^3 \\right\\rangle H_5(\\nu) \\right]\\,. \\label{eq:2D_K3} \\end{equation}  Introducing scaled Hermite polynomials ${\\cal H}_n^{\\pm}(\\nu,\\sigma)\\equiv \\sigma^{\\pm n} H_n\\left(\\nu/\\sigma\\right)$, the polynomial $K_2(\\nu,\\gamma)$, the only one that determines the distribution of saddle points, can be written as \\begin{eqnarray} K_2 =\\frac{1}{8\\pi \\sqrt{3}} \\left[\\ \\vphantom{\\frac{1}{6}} 1  \\right. - \\left( \\left\\langle x q^2 \\right\\rangle + \\frac{1}{3} \\left\\langle x {J_1}^2 \\right\\rangle - \\frac{4}{3} \\left\\langle x J_2 \\right\\rangle + \\frac{2}{3} \\gamma \\left\\langle q^2 J_1 \\right\\rangle + \\frac{2}{9} \\gamma \\left\\langle {J_1}^3 \\right\\rangle - \\frac{2}{3} \\gamma \\left\\langle J_1 J_2 \\right\\rangle \\right) {\\cal H}_1^- \\left(\\nu,\\sqrt{1-2/3\\gamma^2} \\right) \\nonumber \\\\ \\left. + \\frac{1}{6} \\left( \\left\\langle x^3 \\right\\rangle + 2 \\gamma \\left\\langle x^2 J_1 \\right\\rangle + \\frac{4}{3} \\gamma^2 \\left\\langle x {J_1}^2 \\right\\rangle + \\frac{2}{3} \\gamma^2 \\left\\langle x J_2 \\right\\rangle + \\frac{8}{27} \\gamma^3 \\left\\langle {J_1}^3 \\right\\rangle + \\frac{4}{9} \\gamma^3 \\left\\langle J_1 J_2 \\right\\rangle \\right) {\\cal H}_3^- \\left(\\nu,\\sqrt{1-2/3\\gamma^2} \\right) \\right]. \\label{eq:2D_K2} \\end{eqnarray}  The remaining term, $K_3(\\nu,\\gamma)$ is the most complicated one. It is expressed as the expansion in ${\\cal H}_n^+(\\nu,\\sqrt{1-\\gamma^2})$: \\begin{eqnarray} \\lefteqn{K_3 =\\frac{(1-\\gamma^2)}{2 (2 \\pi)^{3/2}(3-2\\gamma^2)^3} \\left[\\vphantom{\\frac{1}{6}} \\gamma (3-2\\gamma^2)^3 {\\cal H}_1^+ \\left(\\nu,\\sqrt{1-\\gamma^2}\\right) + \\left( \\frac{1}{2}\\gamma^3 \\left(1+\\gamma^2-26 \\gamma^4 +28 \\gamma^6 - 8 \\gamma^8 \\right) \\left\\langle x ^3\\right\\rangle \\right.\\right.} \\nonumber \\\\ &&\\left. -\\gamma^4 \\left(26-28 \\gamma^2+8 \\gamma^4\\right) \\left\\langle x^2 J_1 \\right\\rangle + \\gamma \\left(1-\\gamma^2\\right)\\left(1+2 \\gamma ^2\\right) \\left(3-2 \\gamma ^2\\right)^2 \\left\\langle x q^2 \\right \\rangle -\\gamma \\left(24-26 \\gamma^2+8 \\gamma^4 \\right) \\left\\langle x  J_1^2\\right \\rangle \\right. \\nonumber \\\\ &&\\left.\\vphantom{\\frac{1}{6}} +\\gamma\\left(15-23 \\gamma^2 + 8 \\gamma^4 \\right) \\left\\langle x J_2\\right \\rangle + 4 (1-\\gamma^2) \\left(3 -2 \\gamma ^2 \\right)^2 \\left\\langle q^2 J_ 1\\right \\rangle -\\left(10 - 12 \\gamma^2 + 4 \\gamma^4 \\right) \\left\\langle J_1^3\\right \\rangle +6 \\left(1-\\gamma^2\\right) \\left(2-\\gamma^2\\right) \\left\\langle J_1 J_2\\right \\rangle \\right) \\nonumber \\\\ && -\\frac{1}{6} \\left(\\vphantom{\\frac{1}{6}} \\gamma \\left( 27+36 \\gamma^2-224 \\gamma^4+192 \\gamma^6-48 \\gamma^8 \\right) \\left\\langle x^3\\right \\rangle + \\left( 108 - 324 \\gamma^2+ 216 \\gamma^4 - 48\\gamma ^6 \\right) \\left\\langle x^2 J_1\\right \\rangle + 6 \\gamma (3 - 2 \\gamma^2)^3 \\left\\langle x q^2 \\right \\rangle \\right. \\nonumber \\\\ &&\\left. \\left. \\vphantom{\\frac{1}{6}} -36 \\gamma \\left\\langle x J_1^2\\right \\rangle -18 \\gamma \\left\\langle x J_2\\right \\rangle - 8 \\gamma^2 \\left\\langle J_1^3\\right \\rangle - 12 \\gamma^2 \\left\\langle J_1 J_2\\right \\rangle \\right) {\\cal H}_2^+\\left(\\nu,\\sqrt{1-\\gamma^2}\\right) \\right]. \\label{eq:2D_K1} \\end{eqnarray} \\end{widetext} Eqs~(\\ref{eq:difnu1})-(\\ref{eq:difnu2}) (together with the next order expansion available online) are the main theoretical result of this paper. ",
        "conclusions": ""
    },
    "1107/1107.3866_arXiv.txt": {
        "abstract": "The conventional wisdom for the formation of the first hard binary in core collapse is that three-body interactions of single stars form many soft binaries, most of which are quickly destroyed, but eventually one of them survives.  We report on direct N-body simulations to test these ideas, for the first time.  We find that both assumptions are often incorrect: 1) quite a few three-body interactions produce a hard binary from scratch; 2) and in many cases there are more than three bodies directly and simultaneously involved in the production of the first binary.  The main reason for the discrepancies is that the core of a star cluster, at the first deep collapse, contains typically only five or so stars.  Therefore, the homogeneous background assumption, which still would be reasonable for, say, 25 stars, utterly breaks down.  There have been some speculations in this direction, but we demonstrate this result here explicitly, for the first time. ",
        "introduction": "Nowadays, it is recognized that the dynamical evolution of a star cluster is mostly understood as follows \\citep{Heggie03}. Two-body relaxation drives core collapse. The core collapse leads to high stellar density region at the cluster center. Owing to the high stellar density, binaries are formed through three-body encounters \\citep{Aarseth71,Heggie75,Hut85,Goodman93}. These binaries play important roles for post-collapse evolution, triggering re-expansion of the core, as well as gravothermal oscillations \\citep{Bettwieser83, Makino96}.  However, it has not been confirmed whether the binaries are actually formed through three-body encounters, although it has been well investigated how they evolve subsequently, after formation\\citep{Spurzem96}. Indirect evidence has been gathered indicating that more than three stars might be involved \\citep{McMillan88}, but no direct observations have been published of the core conditions that lead to the formation of the first binary, in a simulation in which a cluster of single stars undergoes core collapse. In contrast, later hardening of binaries has been described in considerable detail by \\cite{McMillan90} (MHM90).  In this paper, we have followed the dynamical evolution of star clusters by performing $N$-body simulations, in order to catch the moment of binary formation. Upon spotting the first binary, we have then analyzed in detail how this binary formed, with how many other stars involved, and in what kind of step-by-step process.  The structure of this paper is as follows. In section \\ref{sec:method}, we describe our method for $N$-body simulations and our strategies for catching the moment of binary formation and for analyzing the event. In section \\ref{sec:results}, we show in detail how the binaries are formed. In section \\ref{sec:summary}, we sum up. ",
        "conclusions": "\\label{sec:summary}  In this paper, we have shown that core collapse of star clusters gives rise to the formation of hard binaries via far more complicated processes than have been hitherto considered.  Most importantly, the traditional approximation of binary formation through simultaneous encounters of three bodies offers a vastly oversimplified picture of what really happens.  Typically, more stars are strongly involved, dynamically, in a complex multi-body dance that in no way can be disentangled to lead to a hierarchy with only three stars in the center.  We have found only one exception, among the five cases studied in great detail.  Even in that case, the traditional picture of slow hardening of an initially quite soft binary doesn't apply.  Rather, that three-body process of binary formation produced a rather hard binary very quickly, without subsequent hardening encounters with other stars.  We can point to a single reason for the fact that the two core assumptions (3-body encounters, and gradual hardening) of the traditional scenario for hard binary formation turn out to be unrealistic.  The main reason for the break-down of those assumptions is the very small core at the moment of core collapse.  What we have found is that, typically, the core at that time contains less than 10 stars, and often only 5 or so stars -- so few that a precise measure of the actual number of stars becomes difficult.  Given the discrete nature of core stars, the use of continuum properties like density for a population of only 5 stars is obviously rather problematic.  If core bounce would occur at somewhat larger core sizes, with cores containing twenty to thirty stars or more, the assumptions of three-body encounters and gradual hardening may well have been correct.  Following the path breaking analytical investigation of the role of binaries by \\cite{Heggie75} in the seventies, subsequent analytical work in the eighties seems to suggest that core sizes during core bounce remained large enough to ensure Heggie's assumptions.  For example, from table 2 of \\cite{Goodman84}, equilibrium values of core sizes between $N=1k$ and $N=16k$ range between twenty and sixty stars.  A more detailed analysis in \\cite{Goodman87} gave rise to his equation (IV.8) suggesting a core size at deep bounce large enough for the core to contain dozens of stars.  Our simulations, in contrast, suggest that deep core collapse causes the number of stars in the core to dip below ten in all cases, independent of the number $N$ of stars for the cluster as a whole. During the final stages before the deepest collapse, the remainder of the cluster is effectively frozen, on the local time scale of the most central regions, and even the order of magnitude of $N$ may become irrelevant.  We have demonstrated this for the range $N = 1k$ to $N = 16k$, and we predict a very similar behavior for larger $N$ values, up to $N = 10^6$.  It is an interesting question, why the analytic estimates for the core size at the time of core bounce have overestimated the number of stars in the core $N_{\\rm c,min}$ by almost an order of magnitude. Our results indicate $N_{\\rm c,min} \\sim 5$ while the analytic estimates have predicted results more like $N_{\\rm c,min} \\sim 30$. One possible reason could be the breakdown of the instantaneous approximation for interactions, used in analytic treatments. With the timescale for formation and subsequent hardening of binaries being comparable to the timescale of core contraction, we may need to extend the analytical approach. We are currently looking into these questions in more detail, continuing the investigations reported here.  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.5]{fig01.eps} \\end{center} \\caption{Time evolution of the number of binaries with more than $10 kT$ ($n_{\\rm b}$). Time is scaled by initial half-mass relaxation time, defined in equation (\\ref{eq:trh}).} \\label{fig:nb} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.2]{fig02.eps} \\end{center} \\caption{Time evolution of core radii in the $N=1k$ runs with seed $1$, $2$, and $3$. The core radii in zoom-$0$ panels are calculated at each $1$ $N$-body time unit ($5.9 \\times 10^{-2} t_{\\rm rh}$). Those in the other panels are calculated at each $0.1 t_{\\rm cr,c}$. The width of the time in panels zoom-$1$ is $1 t_{\\rm rh,i}$. The range of the time in panel zoom-$n$ is the same as that between two vertical dashed lines in panel zoom-($n-1$). The ranges between two vertical dashed lines in panels zoom-2 are the same as those of the time in figure~\\ref{fig:sep_n01k_seed01}, \\ref{fig:sep_n01k_seed02}, and \\ref{fig:sep_n01k_seed03}. The arrows and the time $\\tau = 0$ show the time when the first binary is formed in each run.} \\label{fig:rc_n01k} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.2]{fig03.eps} \\end{center} \\caption{Time evolution of core radii in the $N=4k$ run. The core radius in zoom-$0$ panel is calculated at each $1$ $N$-body time unit ($1.8 \\times 10^{-2} t_{\\rm rh}$). Those in the other panels are calculated at each $0.1 t_{\\rm cr,c}$. The width of the time in panels zoom-$1$ is $1 t_{\\rm rh,i}$. From zoom-$1$ to zoom-$3$, the time is zoomed out by ten times step by step. The meanings of the dashed lines are the same as those in figure~\\ref{fig:rc_n01k}. The range between the two vertical dashed lines in panel zoom-$4$ is the same as that of the time in figure~\\ref{fig:sep_n04k_seed01}. The arrows and the time $\\tau = 0$ show the time when the first binary is formed.} \\label{fig:rc_n04k} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.2]{fig04.eps} \\end{center} \\caption{Time evolution of core radii in the $N=16k$ run. The arrows and the time $\\tau = 0$ show the time when the first binary is formed. The core radius in zoom-$0$ panel is calculated at each $1$ $N$-body time unit ($5.4 \\times 10^{-3} t_{\\rm rh}$). Those in panels zoom-$1$, $2$, $3$, and $4$ are calculated at each $1 t_{\\rm cr,c}$, and those in panel zoom-$5$ at each $0.1t_{\\rm cr,c}$. The width of the time in panels zoom-$1$ is $1 t_{\\rm rh,i}$. From zoom-$1$ to zoom-$4$, the time is zoomed out by ten times step by step. The meanings of the dashed lines are the same as those in figure~\\ref{fig:rc_n01k}. The range between the two vertical dashed lines in panel zoom-$4$ is the same as that of the time in figure~\\ref{fig:sep_n16k_seed01}.} \\label{fig:rc_n16k} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.5]{fig05.eps} \\end{center} \\caption{Time evolution of the number of stars in core at around binary formation in the $N=1k$ runs with seed 1, 2, and 3, $N=4k$ run, and $N=16k$ run.} \\label{fig:nc} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.5]{fig06.eps} \\end{center} \\caption{ Orbits of stars involving the binary formation in the $N=1k$ run with seed 1. The stars whose orbits are drawn by red and blue curves finally become the binary components, and their directions are indicated by black arrows. The time ranges are the same as those in panels zoom-$2$, $4$, and $5$ in figure~\\ref{fig:rc_n01k}, \\ref{fig:rc_n04k}, and \\ref{fig:rc_n16k}, respectively.} \\label{fig:xyz} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.35]{fig07.eps} \\end{center} \\caption{Time evolution of separations among stars involving binary formation in the $N=1k$ run with seed 1. The colors are the same as those in figure~\\ref{fig:xyz}. The black curves indicate stars which do not involve the binary formation. For each star, following the color, the number between parentheses refers to the numbers used in the discussion in section \\ref{sec:history}.} \\label{fig:sep_n01k_seed01} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.35]{fig08.eps} \\end{center} \\caption{Time evolution of separations among stars involving binary formation in the $N=1k$ run with seed 2.} \\label{fig:sep_n01k_seed02} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.35]{fig09.eps} \\end{center} \\caption{Time evolution of separations among stars involving binary formation in the $N=1k$ run with seed 3.} \\label{fig:sep_n01k_seed03} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.5]{fig10.eps} \\end{center} \\caption{Time evolution of separations among stars involving binary formation in the $N=4k$ run.} \\label{fig:sep_n04k_seed01} \\end{figure*}  \\begin{figure*} \\begin{center} \\includegraphics[scale=1.35]{fig11.eps} \\end{center} \\caption{Time evolution of separations among stars involving binary formation in the $N=16k$ run.} \\label{fig:sep_n16k_seed01} \\end{figure*}"
    },
    "1107/1107.3920_arXiv.txt": {
        "abstract": "Gravitational flexion, caused by derivatives of the gravitational tidal field, is potentially important for the analysis of the dark-matter distribution in gravitational lenses, such as galaxy clusters or the dark-matter haloes of galaxies. Flexion estimates rely on measurements of galaxy-shape distortions with spin-1 and spin-3 symmetry. We show in this paper that and how such distortions are generally caused not only by the flexion itself, but also by coupling terms of the form (shear $\\times$ flexion), which have hitherto been neglected. Similar coupling terms occur between intrinsic galaxy ellipticities and the flexion. We show, by means of numerical tests, that neglecting these terms can introduce biases of up to 85\\% on the $F$ flexion and 150\\% on the $G$ flexion for galaxies with an intrinsic ellipticity dispersion of $\\sigma_{\\epsilon}=0.3$. In general, this bias depends on the strength of the lensing fields, the ellipticity dispersion, and the concentration of the lensed galaxies. We derive a new set of equations relating the measured spin-1 and spin-3 distortions to the lensing fields up to first order in the shear, the flexion, the product of shear and flexion, and the morphological properties of the galaxy sample. We show that this new description is accurate with a bias $\\leq 7\\%$ (spin-1 distortion) and $\\leq 3\\%$  (spin-3 distortion) even close to points where the flexion approach breaks down due to merging of multiple images. We propose an explanation why a spin-3 signal could not be measured yet and comment on the potential difficulties in using a model-fitting approach to measure the flexion signal. ",
        "introduction": "In the last years weak gravitational lensing became a standard technique to infer the mass distribution of galaxy clusters. In contrast to X-ray or galaxy dynamics, which typically assume hydrostatic or virial equilibrium and spherical mass distributions, gravitational lensing allows to directly probe the projected gravitational potential: Light from distant galaxies gets deflected by the cluster potential and the galaxy ellipticities are modified by the gravitational tidal field of the intervening matter. Hence measurements of the galaxy ellipticities allow to directly trace the curvature of the gravitational potential.  Beyond that, higher-order distortions of background galaxies, like curvature or distortions with three-fold rotational symmetry, can be related to higher-order derivatives of the gravitational potential, thereby providing information on the variation of the potential at smaller scales. By expanding the lens equation to second order, one can identify two combinations of third-order derivatives of the potential, called $F$ and $G$ flexions. These two fields can be related to a skewness in the light distribution and distortions with three-fold rotational symmetry in the shape of galaxies, respectively \\citep{Goldberg05,Bacon06,Irwin07}. Reliable measurements of these higher-order distortions can increase the accuracy with which the inner profile of galaxy clusters can be constrained or can support the detection of substructures \\citep{Bacon06}, whose presence in galaxy clusters is a firm prediction of the Cold-Dark Matter cosmogony \\citep{White78}. Moreover, at galactic scales the incorporation of flexion constraints significantly improves the recovery of halo ellipticity in comparison to shear-only measurements \\citep{Hawken09}.  Two main approaches to measure these higher-order distortions in the galaxy shape have been proposed so far: shapelets \\citep{Goldberg05,Velander11} and HOLICs \\citep{Okura07,Okura08}. Both assume that the measured  distortions can simply be related to the $F$- and $G$-flexions, respectively. This is true only if the shear can be considered negligible and galaxies are intrinsically circular. We will show that a considerable error on the inferred lens power is introduced if any of these assumptions is violated. Then, coupling terms of the form (shear $\\times$ flexion) or (ellipticity $\\times$ flexion) create additional distortions, which are indistinguishable from the pure flexion-induced distortions and thus interfere with the interpretation of the lensing signal.  In \\autoref{sec:flexion_moments} we derive a new set of equations, relating the measured distortion to the lensing fields, including these coupling terms. In \\autoref{sec:morphology} we discuss the physical meaning of the newly identified terms. In \\autoref{sec:new_estimators} we derive a simplified set of equations, applicable to lenses with non-vanishing shear, which also consider the responsivity of the higher-order distortions to shear. These new equations are more general than the equation originally presented by \\citet{Okura08} and constitute a practical, yet accurate approximation to the equations presented by \\citet{Schneider08}.  By means of numerical tests, we show in \\autoref{sec:tests} that our equations are sufficient to reliably describe the lensing-induced change of galaxy shapes up to the point where the very concept of flexion estimates from individual galaxies breaks down due to merging of multiple images of the same source \\citep{Schneider08}. In \\autoref{sec:discussion} we propose an explanation why the $G$-flexion-related spin-3 field has not been measured yet \\citep{Goldberg05,Leonard07, Okura08, Leonard11} despite the fact that the $G$-flexion extends to larger cluster-centric distances than the $F$-flexion, and we comment of the conceptual difficulties of using a model-fitting approach to measure the flexion signal. We conclude in \\autoref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this work we show that measurements of spin-1 ($\\zeta$) and spin-3 distortions ($\\delta$) of galaxy shapes in terms of surface-brightness moments cannot be related in a straightforward way to the $F$- and $G$ flexions caused by gravitational lensing, in particular not for strongly elliptical galaxies or near the centers of galaxy clusters. In the latter case, the reduced shear is not small, thus terms coupling shear and flexion produce non-negligible spin-1 and spin-3 distortions, which contribute to the distortions induced by the flexion alone. A similar effect applies in cases where the observed ellipticity is not small. These coupling terms imply that in general the measured signal will typically be larger than the naively expected flexion signal.  Furthermore, lensing causes a centroid shift that is approximately linear in the flexion fields and it depends on the ellipticity and the concentration of the source galaxy. Hence, when flexion-related distortions measured from individual objects are averaged over a source galaxy population with given distributions of ellipticity and concentration, the measured signal will depend not only on the strength of the lensing fields but also on average galaxy properties.  Our new derivation of the lensing-induced changes of the spin-1 and spin-3 distortion estimators accounts for all of the effects mentioned above. This work extends that of \\cite{Okura08} -- yielding a linear relation between flexion and spin-1 and spin-3 distortions (cf. \\autoref{eq:simpleflex}) -- in two ways: \\begin{itemize} \\item We include shear-flexion coupling terms in the transformation of the flexion estimators $\\zeta$ and $\\delta$; \\item We compute how the average of $\\delta$ and $\\zeta$ over a population of galaxies is affected by the galaxy ellipticity dispersion. \\end{itemize}  We provide two more detailed equations relating $\\langle \\zeta \\rangle$ and $\\langle \\delta \\rangle$ with the reduced shear and the reduced flexion. These equations take the contribution to the measured signal, originating from concentration and ellipticity dispersions of the galaxy population, explicitly into account. Treating the dependence on galaxy morphology is important, since it impacts on the distortion estimates even if shear-flexion cross-terms can be neglected.  By means of several numerical tests, we show that simple linear estimators can lead to biases of up to 85\\% and 150\\% for $\\zeta$ and $\\delta$, respectively. This bias depends in general on the strength of the lensing field, on the intrinsic ellipticity dispersion and on the concentration of the galaxies. Moreover, we showed that the description of the spin-1 and spin-3 distortions that we derived in this paper is valid (with a bias $\\leq 7\\%$ (spin-1 distortion) and $\\leq 3\\%$  (spin-3 distortion)) even close to the point where the flexion concept breaks down due to the merging of multiple images.  We finally find that the measurement of the spin-1 distortion is always stronger (at scales where flexion is usually measured) than the spin-3 distortion. This, combined with the different noise properties of the $\\zeta$ and $\\delta$ estimators, could be an explanation of why a spin-3 field has not been detected so far."
    },
    "1107/1107.2926_arXiv.txt": {
        "abstract": "{ For more than 15 years, since the days of the Energetic Gamma-Ray Experiment Telescope (EGRET) on board the \\textit{Compton Gamma-Ray Observatory} (CGRO; 1991-2000), it has remained an open question why the prominent blazar 3C$\\,$345 was not reliably detected at $\\gamma$-ray energies $\\geq$\\,20\\,MeV. Recently a bright $\\gamma$-ray source (0FGL\\,J1641.4+3939/1FGL\\,J1642.5+3947), potentially associated with 3C$\\,$345, was detected by the Large Area Telescope (LAT) on \\textit{Fermi}. Multiwavelength observations from radio bands to X-rays (mainly GASP-WEBT and \\textit{Swift}) of possible counterparts (3C$\\,$345, NRAO$\\,$512, B3$\\,$1640$+$396) were combined with 20 months of \\textit{Fermi}-LAT monitoring data (August 2008 - April 2010) to associate and identify the dominating $\\gamma$-ray emitting counterpart of 1FGL\\,J1642.5+3947. The source 3C$\\,$345 is identified as the main contributor for this $\\gamma$-ray emitting region. However, after November 2009 (15 months), a significant excess of photons from the nearby quasar NRAO$\\,$512 started to contribute and thereafter was detected with increasing $\\gamma$-ray activity, possibly adding flux to 1FGL\\,J1642.5+3947. For the same time period and during the summer of 2010, an increase of radio, optical and X-ray activity of NRAO$\\,$512 was observed. No $\\gamma$-ray emission from B3$\\,$1640$+$396 was detected. } ",
        "introduction": "The quasar \\object{3C$\\,$345} is known as a prominent variable source from radio- to X-ray bands ($\\leq$\\,10\\,keV). It is particularly bright at radio wavelengths and has an extended radio structure that is observable from sub-pc to kpc scales which is archetypical for a relativistic blazar jet \\citep[e.g.][]{1999ApJ...521..509L, 1989AJ.....97.1550K}. Its jet has an apparent opening angle of $\\sim$\\,13\\degr\\ \\citep{2009A&A...507L..33P}. \\citet{1994ApJ...432..103U} concluded that the observed high-energy emission of this source is caused by inverse Compton up-scattering of lower energy photons by relativistic electrons of the pc-scale radio jet. However, it remained a puzzle why 3C$\\,$345 was not detected at energies above 20\\,MeV, while 3C$\\,$279 with similar apparent properties is a prominent source at $\\gamma$-ray energies.  The high-redshift quasar \\object{NRAO$\\,$512}, 0.5\\degr\\ west of 3C$\\,$345, is known as a compact radio source with a two-sided jet morphology on kpc scales and with a flat radio spectrum \\citep{1999yCat..41390545K}. The observed VLBI radio structure revealed a wide apparent jet opening angle of $\\sim$\\,50\\degr\\ \\citep{2009A&A...507L..33P}, which implies a small viewing angle to the line of sight, with its emission directly beamed at us.  The Energetic Gamma-Ray Experiment Telescope \\citep[EGRET;][]{1993ApJS...86..629T} on board the \\textit{Compton Gamma Ray Observatory} (CGRO) spacecraft did not observe a significant excess of $\\gamma$-ray photons consistent with 3C$\\,$345 \\citep[3EG;][]{1999ApJS..123...79H}, covering the first four phases of the CGRO mission (1991-1995). This lack of detectable flux of $\\gamma$-rays above 20\\,MeV from 3C$\\,$345 was attributed to the high detection threshold of EGRET. A source consistent with the position of 3C$\\,$345, labeled as EGR$\\,$J1642+3940, was found by \\citet{2008A&A...489..849C}. This source was seen almost entirely in a single two-week viewing period with an excess of $\\sim$\\,5.8$\\sigma$, from April 23rd - May 7th, 1996 (vp5190), which was after the data used for the 3EG. Unfortunately, the angular resolution of EGRET was too poor to distinguish contributions from the other two nearby candidates NRAO$\\,$512 and B3$\\,$1640$+$396, see Table~\\ref{tab:sources}. Between October 1st and October 4th, 1993, a similar source was marginally detected by EGRET ($\\sim$\\,2.1$\\sigma$), but its centroid was localized closer to Mrk$\\,$501. In exactly the same viewing period, \\citet{1999ApJ...514..138K} claimed EGRET detected a flare from Mrk$\\,$501, which lies 2.5\\degr\\ east of 3C\\,345. This would then be the only time that either of those sources was seen by EGRET -- a strange coincidence that led Casandjian \\& Grenier to note that the EGR$\\,$J1642+3940 source association with 3C$\\,$345 is ``not clear''.  On June 11th, 2008, the \\textit{Fermi} Gamma-ray Space Telescope (short \\textit{Fermi}) was launched, equipped with the Large Area Telescope (LAT) instrument \\citep{2009ApJ...697.1071A} providing continuous all-sky monitoring of $\\gamma$-ray emission as a default mode of operation. The LAT, compared with EGRET, has an improved sensitivity by a factor of $\\sim$\\,25 and an improved point-source localization. After 20 months of LAT operation within the region of interest, a localization improvement (comparing radii of 95\\% error circles) of a factor of $\\sim$\\,14 was reached compared to EGRET. These improvements provided a new opportunity to probe 3C$\\,$345 for $\\gamma$-ray emission above 100\\,MeV.  The three-month bright \\textit{Fermi} source list (0FGL) for the dataset collected between August and October 2008 reported a $\\gamma$-ray source (0FGL\\,J1641.4+3939) near the position of 3C$\\,$345. It was then associated with the quasar B3$\\,$1640$+$396 (CLASS$\\,$J1641+3935) at low confidence \\citep{2009ApJ...700..597A}. The bright 0FGL source was listed in the first \\textit{Fermi}-LAT source catalog (1FGL) as 1FGL\\,J1642.5+3947 \\citep{2010ApJS..188..405A}, further constraining its localization. An association with 3C$\\,$345 was still not possible with high confidence based on \\textit{Fermi}-LAT data alone, nevertheless, 1FGL\\,J1642.5+3947 was listed in the catalog as being affiliated with its most likely counterpart 3C$\\,$345 \\citep{2010ApJ...715..429A}.  In October 2009 a GeV flare was detected from the vicinity of 0FGL\\,J1641.4+3939 \\citep{2009ATel.2226....1R}. Simultaneously, the GASP-WEBT collaboration reported on increased optical and radio activity of 3C$\\,$345 \\citep{2009ATel.2222....1L}. This event triggered additional observations at different frequencies, which made it possible to establish an identification through correlated multiwavelength activity.  Here we present 20 months of monitoring observations by \\textit{Fermi} LAT of the region around 3C$\\,$345, NRAO$\\,$512 and B3$\\,$1640$+$396, as well as radio, optical, UV, and X-ray observations of these sources. Section 2 discusses the observations and data reduction methods applied to obtain calibrated datasets for the subsequent analysis. Section 3 discusses the results of \\textit{Fermi}-LAT observations and identification through multiwavelength counterparts of events in the light curve. Section 4 discusses and summarizes the findings of this paper. ",
        "conclusions": "Properties of the $\\gamma$-ray emitting region reported in the first three-month Bright Source List \\citep{2009ApJS..183...46A} as well as in the first 11-month \\textit{Fermi}-LAT source catalog \\citep{2010ApJS..188..405A, 2010ApJ...715..429A} were compared to the 20-month dataset presented here. Localizations of the sources in this region were improved over the previously reported values, leading to the detection of two $\\gamma$-ray point sources, one consistent with 1FGL\\,J1642.5+3947 and a new $\\gamma$-ray source $\\sim$\\,0.4\\degr\\ west of it. 3C$\\,$345 and NRAO$\\,$512 were identified as the counterparts for these $\\gamma$-ray sources. Their $\\gamma$-ray spectra have a spectral slope of $\\Gamma=$\\,2.4--2.5 typical for that of flat spectrum radio quasars \\citep{2010ApJ...715..429A}. No significant spectral break was observed in the spectrum of either source, unlike in the case of 3C 454.3 \\citep{2010ApJ...721.1383A}, however, the statistics at the highest energies are poor, and the presence of a spectral break at E$>$10\\,GeV is not ruled out.  The $\\gamma$-ray emission from 3C$\\,$345 was identified based on an improved $\\gamma$-ray counterpart localization and multi-wavelength activity including correlations of moderate significance found between the optical and radio variability, using a plausible range of $\\beta$ in the $\\gamma$-ray and optical PSDs, and major $\\gamma$-ray events observed by \\textit{Fermi} LAT during its first 20 months of operation. EGRET observed 3C$\\,$345 between 1991 and 1994 during a time of enhanced radio activity. An upper limit of 2.5\\,$\\cdot$\\,10$^{-7}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ for the viewing period September 12th - 19th, 1991, was reported \\citep{1994ApJS...94..551F}. However, \\citet{2008A&A...489..849C} report the detection of a $\\gamma$-ray source that could be consistent with 3C\\,345 in the period April 23rd - May 7th, 1996 with a $\\gamma$-ray flux of (3.5\\,$\\pm$\\,0.8)\\,$\\cdot$\\,10$^{-7}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$, but it was noted that the association with 3C$\\,$345 remains unclear because of a possible confusion with Mrk\\,501. In the \\textit{Fermi}-era, with an increased sensitivity and the possibility of long-term integrations, 3C$\\,$345 was already detected during the first three months of operation at a level of $\\sim$\\,18$\\sigma$ and after 20 months it reached over 30$\\sigma$. However, owing to the unclear association of 3C\\,345 with EGR\\,J1642+3940, the three-month data were falsely associated with CLASS\\,J1641+3035 (B3\\,1640+396), which, as shown in this paper, did not produce a significant excess of $\\gamma$-ray emission above the detection threshold of \\textit{Fermi}-LAT. Thus the simplest explanation for the previous non-detection of 3C$\\,$345 is that its emission was just below or at the detection limit of EGRET.  These findings end the decade-long debate whether 3C$\\,$345 is $\\gamma$-ray loud \\citep{1997ApJ...480..596U}. The identification of 1FGL\\,J1642.5+3947 with 3C$\\,$345 is further supported by the connection of radio-emission of the parsec-scale jet with the observed $\\gamma$-ray emission \\citep{SchinzelFmJ2010}. Here not only individual flares were correlated with radio events, but an unambigious long-term trend in $\\gamma$-ray emission was observed, which matched the trend of the radio flux density light curve of the inner jet at 43\\,GHz. The trend and correlated radio events relate to a distance of $\\sim$10$\\,$pc in the source frame and questions existing scenarios that place the $\\gamma$-ray emission at sub-pc distances from the central engine \\citep[e.g.,][]{2010MNRAS.405L..94T}. This supports the arguments given by \\citet{1997ApJ...480..596U} that the observed $\\gamma$-ray emission is related to the relativistic jet on pc scales.  The cross-correlation found between $\\gamma$-ray and optical light curves at 110\\,days had a moderate confidence level of $>$\\,98.5\\%, whereas the lag at -20\\,days had a confidence level of only $>$\\,67.6\\% using a plausible range of $\\beta$ in the $\\gamma$-ray and optical PSDs. If typical values for the power spectral density (PSD) shapes are assumed ($\\beta_{\\gamma} = 1.5$ and $\\beta_{\\mathrm{optical}} = 1.5$; second to lowest dashed line in the top panel of Fig.~\\ref{fig:dcf}), the confidence level remains low at $>$\\,89.3\\%. The weak correlation found between the $\\gamma$-ray and 230\\,GHz radio light curves had a worst-case confidence level of $>$\\,89.3\\%. Assuming a possibly more typical combination of PSD shapes ($\\beta_{\\gamma} = 1.5$ and $\\beta_{\\mathrm{radio}} = 2.0$), a confidence level of $>$\\,96.1\\% is obtained. If this weak correlation holds, it could suggest a connection between the radio emission and production of $\\gamma$-rays. A timelag of $\\sim$-120 days suggests the mm-radio emission precedes the $\\gamma$-rays. \\citet{2010ApJ...722L...7P} showed for a larger sample of radio-loud quasars that on average the 15\\,GHz radio emission lags $\\sim$1.2 months behind $\\gamma$-rays. The result presented here indicates that between 15\\,GHz and 230\\,GHz it is possible to directly observe the $\\gamma$-ray emitting region, as already implied by \\citet{1997ApJ...480..596U}, which supports the argument of $\\gamma$-ray production in the relativistic jet on pc scales. However, in contrast to these, a more confident cross-correlation result of $\\gamma$-rays leading by $\\sim$80 and $\\sim$50\\,days was recently obtained for two prominent BL Lac type objects, OJ\\,287 \\citep{2011ApJ...726L..13A} and AO\\,0235+164 \\citep{2011arXiv1105.0549A}, respectively. Doubts on the indicated radio/$\\gamma$-ray correlation of 3C\\,345 remain in this problem and should be investigated more thoroughly with a longer and denser sampled lightcurve at mm wavelengths.  The $\\gamma$-ray emission from NRAO$\\,$512, a high-redshift quasar, was identified based on a correlated increase in multiwavelength emission from radio up to GeV energies between 2009 and 2010. At monthly timescales the source started to be significant around November 2009. This increase of $\\gamma$-ray flux led to the first detection of the now identified new $\\gamma$-ray counterpart west of 3C$\\,$345. The LAT flux quoted in this paper for the first 15 months was below the EGRET detection threshold, offering a plausible explanation of why EGRET did not see emission from NRAO$\\,$512, except possibly during a $\\gamma$-ray flare event.  The radio, optical and X-ray flux ratios between 3C$\\,$345, NRAO$\\,$512 and B3$\\,$1640$+$396 could be used to gauge their individual $\\gamma$-ray photon contributions and might be able to provide an indication of false associations for $\\gamma$-ray photons from this part of the sky. Multiwavelength monitoring, in addition to the continuous all-sky monitoring by \\textit{Fermi} LAT, will help to improve the localization of $\\gamma$-ray emission from this region and is necessary for a deeper variability study for any of these sources in the future. Ultimately this might lead to the detection of $\\gamma$-ray emission from B3$\\,$1640$+$396 as well."
    },
    "1107/1107.5737_arXiv.txt": {
        "abstract": "Recent detections of GeV photons in a few gamma-ray bursts (hereafter GRBs) by \\textit{Fermi}-LAT imply huge bulk Lorentz factors to avoid a large $\\gamma\\gamma$ optical depth at high energy. Estimates can be as high as $\\Gamma\\simeq 10^3$ in the most extreme cases. This puts severe constraints on models of the central engine and the jet acceleration in GRBs. These estimates are however obtained from a simplified single zone model. We present here a more realistic calculation of the $\\gamma\\gamma$ opacity which takes into account the time, space and direction dependent photon field existing in an outflow with several relativistically moving emitting zones. The formalism is very general and can be applied to many models of the prompt GRB emission. We present results obtained for a numerical implementation in the framework of the internal shock model. We show that (i) the minimum Lorentz factor $\\Gamma_\\mathrm{min}$ in bright \\textit{Fermi}-LAT GRBs is reduced by a factor $\\simeq 2-3$ compared to previous estimates if the GeV and MeV emission are produced in the same region, and by an additional factor $\\simeq 2-8$ if the GeV emission is produced at larger radii. We provide an improved approximate formula for $\\Gamma_\\mathrm{min}$ which is in good agreement with our numerical results and which can be directly applied to \\textit{Fermi}-LAT GRB data; (ii) a delayed onset of the GeV emission can be due to the time evolution of the opacity in a GRB outflow. As an illustration of these first two results, we present a synthetic GRB that reproduces most features of GRB 080916C with a mean Lorentz factor of $\\simeq 340$, an optically thin regime for $\\gamma\\gamma$ opacity at 3 GeV in time bin 'b', a variability timescale of $\\simeq 0.5$ s in the MeV lightcurve and a delayed onset of $\\simeq 5$ s of the GeV emission; (iii) the $\\gamma\\gamma$ opacity can smooth the short timescale variability in the GeV lightcurve. This last result implies that the observed variability at high energy is not necessarily a good test to distinguish between an internal and an external origin for the GeV emission in GRBs. ",
        "introduction": "The combination of  a short timescale variability with a huge radiated gamma-ray energies in gamma-ray bursts leads to the well-known compactness problem: $\\gamma$-ray photons should not escape due to $\\gamma\\gamma$ annihilation. This contradiction is solved by assuming that the emitting region has an ultra-relativistic motion \\citep{rees:1966}. This leads to an important constraint on the minimum Lorentz factor of the outflow in cases where the GRB spectrum does not show any evidence for a high-energy cutoff \\citep[see e.g.][]{baring:1997,lithwick:2001}, or to a direct measurement of the Lorentz factor if such a cutoff is identified, as possibly in GRB 090926A \\citep{ackermann:2011}.\\\\    Since the launch of \\textit{Fermi} in June 2008, the LAT instrument has detected high energy photons above 10 GeV in a few GRBs. The observed $\\gamma$-ray spectrum often remains consistent with a Band function covering the GBM and LAT spectral ranges without any evidence of a high energy cutoff which could be identified as a signature of $\\gamma\\gamma$ annihilation (with possibly the exception of GRB 090926A, \\citealt{ackermann:2011}). This extension by \\textit{Fermi} of the observed spectral range upper bound from a few MeV to a few tens of GeV implies constraints on $\\Gamma_{\\mathrm{min}}$ which are much more severe than the ones obtained previously \\citep[see e.g.][]{racusin:2011}. In a few cases $\\Gamma_{\\mathrm{min}}$ has been estimated to be of the order of 1000. For instance, in the four brightest \\textit{Fermi}-LAT GRBs, the estimated minimum values are $\\Gamma_{\\mathrm{min}}\\simeq 887$ in GRB 080916C \\citep{abdo:2009a},  $\\simeq 10^3$ in GRB 090902B \\citep{abdo:2009b}, $\\simeq 1200$ in GRB 090510  \\citep{ackermann:2010}), $\\simeq 720$ in GRB 090926A \\citep{ackermann:2011}. These extreme values put severe constraints on the physics of the central engine which should be able to strongly limit the baryon load in the outflow. However these minimum Lorentz factors were obtained from a simplified ``single zone'' model where the spatial and temporal dependencies of the radiation field are averaged out. This calls for more realistic calculations of the $\\gamma\\gamma$ opacity in GRB outflows. We present here a detailed approach taking into account the exact time, space and direction dependent radiation field in outflows with several relativistically moving emitting regions, such as internal shocks. \\\\  The paper is organized as follows. In \\refsec{section_calculation_tgg}, we derive the exact formula for the $\\gamma\\gamma$ opacity in complex and variable GRB outflows, we describe its dependency on several physical parameters such as the emission radius or the Lorentz factor and we compare the exact calculation to the approximate formula which is broadly used in the literature. Finally we validate our approach by a comparison with the semi-analytical study of \\citet{granot:2008}. \\refsec{section_internal_shocks} is devoted to the investigation of several consequences of $\\gamma\\gamma$ annihilation in GRBs in the framework of the internal shock model. We discuss successively the constraints on the minimum Lorentz factor, with a focus on GRB 080916C, the spectral shape of the high-energy cutoff, the delayed onset of the GeV emission and the smoothing of the short timescale variability in GeV lightcurves. We compare also in this section the  $\\gamma\\gamma$ opacity with other expected sources of opacities. \\refsec{sec_2zones_model} presents a discussion of models where the GeV emission is produced at a larger radius than the MeV main component. Finally, \\refsec{sec_conclusions} is the conclusion. ",
        "conclusions": "\\label{sec_conclusions} Photon--photon annihilation is an important process in GRBs. It has been known for a long time that it provides one of the main constraints on GRB models, i.e. the minimum Lorentz factor in the jet. In most GRBs, this leads to $\\Gamma_\\mathrm{min}\\simeq 10^{2}$ \\citep{lithwick:2001}. The detection of GeV photons in a few GRBs by \\textit{Fermi}-LAT has led to much stronger limits, of the order of $\\Gamma_\\mathrm{min}\\simeq 10^3$ \\citep{abdo:2009a}. However these estimates are based on highly simplified single zone formulae, which assume an isotropic radiation field in the coming frame of the outflow. We present here a detailed calculation of the $\\gamma\\gamma$ opacity in GRB outflows taking into account the exact time-dependent anisotropic photon field produced by multiple relativistically geometrically thin emitting regions. This calculation can be implemented numerically, for any model prescription of the prompt emission mechanism, as long as the emitting regions can be considered as expanding thin spherical shells. We present results obtained in the framework of the internal shock model, where the emitting regions are shock waves propagating within the outflow. Our calculation has been validated by a comparison with the previous semi-analytical study made by \\citet{granot:2008}. We have also estimated the opacity due to Thomson scatterings by ambient electrons in the outflow, to pairs created by $\\gamma\\gamma$ annihilation and to $\\gamma\\gamma$ annihilation with prompt photons back-scattered by the external medium. For all the examples presented in the paper, the $\\gamma\\gamma$ opacity is the dominant term (see \\refpar{sec_opacity_others}). We have obtained the following results: \\begin{enumerate} \\item \\textit{Spectral shape of the $\\gamma\\gamma$ cutoff:} as described by \\citet{granot:2008}, due to a time evolution of the $\\gamma\\gamma$ opacity, the spectral shape of the $\\gamma\\gamma$ cutoff in time-integrated spectra is not expected to be exponential, but rather power-law like. We illustrate this effect in a synthetic burst in \\refpar{subsection_scattering}. This spectral shape makes more difficult the identification of the cutoff in \\textit{Fermi}-LAT spectra. The main candidate is GRB 090926A \\citep{ackermann:2011}, where the number of photons detected above the cutoff at $\\simeq 1.4$ GeV is too low to allow for a precise characterization of the spectral shape. Nevertheless, in most \\textit{Fermi}-LAT GRBs, there is no evidence for an observed cutoff, which puts interesting constraints on the outflow.\\\\  \\item \\textit{Minimum Lorentz factor in GRB outflows:} when taking into account the exact photon field, the $\\gamma\\gamma$ opacity for GeV photons is reduced compared to single zone isotropic estimates, leading to minimum Lorentz factors which are a factor of $\\simeq2-3$ lower. As an illustration, we present in \\refpar{subsection_lfmin} a synthetic burst that reproduces well the observed features of time bins 'a' and 'b' of GRB 080916C with a mean Lorentz factor of $\\simeq 340$. This minimum Lorentz factor is obtained assuming that there is only one component in the emitted spectrum, responsible both for the MeV and the GeV photons. This is not necessarily the case as several \\textit{Fermi}-LAT burst show evidence of an additional high-energy component. We show in \\refpar{sec_gmin_2zones} that the minimum Lorentz factor can be reduced by another factor $2$ to $8$ if the GeV photons are emitted at larger radius than the main component. This study clearly illustrates the need for a detailed modeling to constrain the Lorentz factor in GRB outflows. When it is not possible, a reasonably accurate estimate of $\\Gamma_\\mathrm{min}$ can be obtained from the following formula \\begin{eqnarray} \\Gamma_\\mathrm{min} & \\simeq &  \\frac{\\left[C_1 2^{1+2\\beta} \\mathcal{I}(\\beta)\\right]^{\\frac{1}{2(1-\\beta)}}}{\\left[\\frac{1}{2}\\left(1+\\frac{R_\\mathrm{GeV}}{R_\\mathrm{MeV}}\\right)\\left(\\frac{R_\\mathrm{GeV}}{R_\\mathrm{MeV}}\\right)\\right]^{1/2}}\\, \\left(1+z\\right)^{-\\frac{1+\\beta}{1-\\beta}} \\nonumber\\\\ % & &\\!\\!\\!\\! \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! \\times % \\left[ \\sigma_\\mathrm{T} \\left(\\frac{D_\\mathrm{L}(z)}{c \\Delta t_\\mathrm{var}}\\right)^2 \\!\\!\\! E_\\mathrm{c} F(E_\\mathrm{c})\\right]^{\\frac{1}{2(1-\\beta)}} \\left(\\frac{E_\\mathrm{max}E_\\mathrm{c}}{(m_\\mathrm{e}c^2)^2}\\right)^{\\frac{\\beta+1}{2(\\beta-1)}} \\!\\!\\!\\! \\, ,\\nonumber\\\\ \\label{eq_newapprox} \\end{eqnarray} where $C_1 \\simeq 4\\cdot 10^{-2}$, $z$ and $D_\\mathrm{L}(z)$ are the redshift and the luminosity distance of the source, $\\Delta t_\\mathrm{var}$ is the observed variability timescale, $R_\\mathrm{GeV}/R_\\mathrm{MeV}$ is the ratio of the radii where the GeV and MeV components are emitted, and where the high energy spectrum (over a duration $\\sim \\Delta t_\\mathrm{var}$) is assumed to follow a power-law with photon index $\\beta$ above an observed characteristic energy $E_\\mathrm{c}$ : $\\overline{F}(E)=\\overline{F}(E_\\mathrm{c})(E/E_\\mathrm{c})^\\beta$ ($\\mathrm{ph.cm^{-2}.keV^{-1}}$). Energy $E_\\mathrm{max}$ is the observed energy of the most  energetic detected photons. The integral $\\mathcal{I}(\\beta)$ is defined by \\refeq{eqn_i_beta} in \\refpar{sec:flash}. As usually the spectrum is measured over a time interval $\\Delta t_\\mathrm{spec}$ which is larger than the variability timescale $\\Delta t_\\mathrm{var}$, the normalization $F(E_\\mathrm{c})$ entering in \\refeq{eq_newapprox} (fluence at energy $E_\\mathrm{c}$ in $\\mathrm{ph.cm^{-2}.keV^{-1}}$) must be corrected by a factor $F(E_\\mathrm{c})=\\overline{F}(E_\\mathrm{c})\\times \\left(\\Delta t_\\mathrm{var}/\\Delta t_\\mathrm{spec}\\right)$. This equation can be directly applied to \\textit{Fermi}-LAT observations and generalizes the usual formula given by \\citet{abdo:2009a} by introducing two corrections: (1) a more accurate normalization taking into account the anisotropy of the radiation field and including a numerical factor $C_1$ obtained from the comparison with numerical simulations in \\refpar{subsection_lfmin}; (2) the possibility to take into account two different emitting regions for MeV and GeV photons. The standard limit is obtained with $R_\\mathrm{GeV}/R_\\mathrm{MeV}=1$ (same region) : then the denominator in \\refeq{eq_newapprox} equals $1$. The radius $R_\\mathrm{MeV}$ is estimated from the variability timescale by $R_\\mathrm{MeV}\\simeq \\Gamma^2 c\\Delta t_\\mathrm{var}/(1+z)$, which is valid for most models of the prompt emission. The radius $R_\\mathrm{GeV}$ is difficult to constrain without a detailed model of the high-energy emission mechanism. If GeV photons have an internal origin, an upper limit for $R_\\mathrm{GeV}$ is given by the deceleration radius. In the future, a measurement of the variability timescale in the GeV lightcurve could provide a better estimate of this radius.\\\\  \\item \\textit{Delayed onset of the GeV emission:} due to the variable nature of the GRB phenomenon, the $\\gamma\\gamma$ opacity is expected to be strongly time dependent. This can lead naturally to a delayed onset of the high-energy component, if it is initially highly absorbed. For instance, in internal shocks, the emission radius increases during the propagation of a given shock wave, which favors such an evolution. We show several examples of synthetic bursts where a delayed onset of the GeV component is observed. In the case of internal shocks, the delay before the onset is comparable with the observed duration associated with the propagation of the shock wave that produces the seed photons for $\\gamma\\gamma$ annihilation. This duration has to be distinguished from the shortest timescale of variability. In the synthetic burst which models time bins 'a' and 'b' of GRB 080916C, we obtain a GeV onset delayed by $\\simeq 5\\, \\mathrm{s}$, comparable with the observed value, while reproducing variability on shorter timescales ($\\simeq 0.5$ s) in the MeV lightcurve. A time evolving $\\gamma\\gamma$ opacity appears as a good candidate to explain the observed delayed onset of the GeV emission. Note that this effect is obtained here assuming a single component in the emitted spectrum. More complex spectral evolution, such as a varying inverse Compton component \\citep{bosnjak:2009}, may be an additional source of delay.\\\\  \\item \\textit{Smoothing the variability in the GeV lightcurve:} in highly variable outflows, the short timescale variability in the GRB lightcurve is expected to be associated with emission regions at low radius. Therefore, the associated high-energy emission may be highly absorbed by $\\gamma\\gamma$ annihilation due to a denser photon field. This predicts a smoothing of the lightcurve in the GeV range. We illustrate this effect in the framework of the internal shock model in \\refpar{sec_smoothing}. For this reason, the variability of the GeV lightcurve is an indicator to use carefully if one wants to distinguish between an internal and an external origin for the GeV emission. \\end{enumerate} Our results loosen the constraints on models of  GRB central engines as Lorentz factor of $10^3$ do not seem to be required, even in the most extreme bursts observed by \\textit{Fermi}-LAT. In addition, they solve another contradiction. For Lorentz factors of $10^3$, the deceleration by the external medium occurs very early. Then, it becomes difficult to interpret \\textit{Fermi}-LAT GRBs such as GRB 080916C if the prompt emission is radiated well above the photosphere (see also \\citealt{zhang:2009}). On the other hand, the early steep decay observed in X-ray afterglows by Swift-XRT is a strong evidence that the prompt emission phase ends at a large radius \\citep{lyutikov:2006,lazzati:2006,kumar:2007,genet:2009}. With lower limits on the Lorentz factor as obtained in the present paper, bright \\textit{Fermi}-LAT GRBs such as GRB 080916C are consistent with models, including internal shocks, where the prompt emission is produced between the photospheric and the deceleration radii.    \\appendix"
    },
    "1107/1107.1256.txt": {
        "abstract": "{We apply the method of Burnett \\& Binney (2010) for the determination of stellar distances and parameters to the internal catalogue of the Radial Velocity Experiment (Steinmetz et al.\\ 2006). Subsamples of stars that either have Hipparcos parallaxes or belong to well-studied clusters, inspire confidence in the formal errors. Distances to dwarfs cooler than $\\sim6000\\,$K appear to be unbiased, but those to hotter dwarfs tend to be too small by $\\sim10\\%$ of the formal errors. Distances to giants tend to be too large by about the same amount. The median distance error in the whole sample of 216\\,000 stars is $28\\%$ and the error distribution is similar for both giants and dwarfs.  Roughly half the stars in the RAVE survey are giants. The giant fraction is largest at low latitudes and in directions towards the Galactic Centre.  Near the plane the metallicity distribution is remarkably narrow and centred on $\\mh=-0.04\\,$dex; with increasing $|z|$ it broadens out and its median moves to $\\mh\\simeq-0.5$.  Mean age as a function of distance from the Galactic centre and distance $|z|$ from the Galactic plane shows the anticipated increase in mean age with $|z|$. } ",
        "introduction": "In recent years there have been significant advances in our knowledge of the Galaxy due to a number of large-scale surveys in complementary magnitude ranges. The Hipparcos mission (\\citealt{Hipparcos}) obtained astrometry of unprecedented precision for a sample of bright stars that was complete only down to $V\\sim8$, and at the other end of the scale we have the Sloan Digital Sky Survey (SDSS, \\citealt{SDSS}), going no brighter than around magnitude $r=14$ and significantly fainter than $r=18$ (\\citealt{Ivezic01}). The sizable gap between these two studies is filled by the Radial Velocity Experiment (RAVE, \\citealt{RaveDR1}), focusing primarily on magnitudes in the range $9<I<13$. RAVE is still ongoing, and has provided a number of important results regarding the structure and kinematics of the Galaxy (see for example \\citealt{Smith07}, \\citealt{Seabroke08}, \\citealt{Munari09}, \\citealt{RAVE_velell}, \\citealt{Arnaud10}).  One of the challenges thrown up by recent surveys, concentrating as they do on relatively distant stars, is that of distance estimation. The astrometric precision required to measure trigonometric parallaxes to these stars is not yet available, so we must rely on secondary methods for determining stellar distances. A useful step in this direction was taken by \\cite{Breddels}, who showed that RAVE data could be used to estimate the parameters and thus distances for objects in the second RAVE data release. Their technique was honed by \\cite{Zwitter10}, who used spectrophotometric distances to characterise the RAVE survey.  An alternative approach was developed by \\cite{burnett1} in which our prior knowledge of the Galaxy and stars is more systematically exploited. In this paper we apply this technique to $\\sim216\\,000$ stars observed by RAVE.  We determine spectrophotometric distances in parallel with other stellar parameters, in particular metallicity, age and mass, so we are able to study how the giant/dwarf ratio and the distributions in age and metallicity vary in the region surveyed by RAVE.  Previous applications of Bayesian inference to the determination of stellar parameters have focused on single parameters such as age \\citep{Jorgensen}; to our knowledge, this is the first study to consider all parameters together.  In Section \\ref{sec:input} we specify the data on which our distances are based, and in Section \\ref{sec:theory} we summarise the algorithm used to calculate distances. Section \\ref{sec:tests} tests the derived distances by comparing them with (i) Hipparcos parallaxes and (ii) established cluster distances. The comparison with Hipparcos parallaxes uncovers slight biases in the distances to hot dwarfs and giants, respectively. Section \\ref{sec:tests} tests the validity of the formal errors by (i) comparing distances to the same stars derived from independent spectra, and (ii) comparing our distances with those obtained by \\cite{Zwitter10}. The algorithm determines probability distributions for several stellar parameters in addition to distance. In Section \\ref{sec:output} we display the distribution of formal errors in metallicity, age and initial mass, in addition to those in distance. In Section \\ref{sec:SF} we use our distances to explore which regions of the Galaxy are probed by the RAVE survey, and indicate which regions are predominantly probed with dwarfs or giants. In Section \\ref{sec:param} whether examine how how mean age and metallicity vary within the probed portion of the Galaxy. Section \\ref{sec:conclusions} sums up. ",
        "conclusions": "\\label{sec:conclusions}  We have derived distances (or parallaxes) to $\\sim216\\,000$ stars in the RAVE survey using the Bayesian analysis of \\cite{burnett1}. We have checked the parallaxes and their associated errors against Hipparcos parallaxes, and conclude that for dwarfs cooler than $\\Teff=6000\\K$ the parallaxes are unbiased, but the parallaxes of hotter dwarfs are systematically too large by $\\sim0.1\\sigma$ because even the VDR2 pipeline over-estimates the gravities of hot dwarfs. The parallaxes of giants tend to be too small by $\\sim0.1\\sigma$. For all three classes of star, hot dwarfs, cool dwarfs and giants, the scatter in differences between the spectrophotometric and Hipparcos parallaxes is consistent with the formal errors on the parallaxes.  We have checked our distances and our errors against distances to star clusters, which tend to be beyond the reach of Hipparcos. This rather small sample of stars, which contains both dwarfs and giants, is consistent with our distances being unbiased and their errors being accurate.  RAVE has obtained more than one spectrum for $\\sim19\\,000$ stars. We have used this sample to assess our errors by comparing independent distances to the same object. The scatter in the difference of distances is only half the quadrature-sum of their formal errors. This is consistent with a significant contribution to the errors coming from factors, such as defects in the spectral analysis and the physics of the stellar models, that are the same for repeated determinations of the distance to a given star.  Comparison of our distances to the spectrophotometric distances of \\cite{Zwitter10} shows that we assign slightly larger distances to dwarfs and smaller distances to giants than do Zwitter et al. Given the signs of our deviations from Hipparcos parallaxes, it follows that for both dwarfs and giants our distances are more accurate than those of Zwitter et al. For dwarfs this finding can be traced to the use by Zwitter et al.\\ of the VDR3 pipeline, which tends to over-estimate $\\log g$, especially for hot dwarfs. Our distances to giants benefit from a more sophisticated prior, which tends to pull stars to smaller distances. For dwarfs the distribution of normalised residuals between our distances and those of Zwitter et al.\\ is rather narrow, having a dispersion of only $0.52$, implying that much of the uncertainty in distance arises from errors in the original data and the spectral template library, which are common to the two studies.  Our formal errors are based on the conservative estimates of the errors in the input data given by \\cite{RaveDR2}. The conservative nature of these errors is confirmed by the analyses of both repeat observations and the distances of \\cite{Zwitter10}. However, the analysis of Hipparcos stars indicates that by happy chance our formal error budget is just large enough to encompass external sources of error, such as spectral mis-match and deficiencies in the stellar models, so our formal errors are close to the final uncertainties in our distances.  We have examined the distributions of the errors returned by the Bayesian analysis for distance, metallicity, age and initial mass. The median formal distance error is $28\\%$, the median formal uncertainty in [M/H] is 0.17\\,dex, the median formal uncertainty in $\\log(\\tau)$ is 0.27\\,dex and the median fractional uncertainty in initial mass is $16\\%$. These figures show that by using prior knowledge of the structure of our Galaxy and the nature of stellar evolution, one can constrain stellar parameters more narrowly than when each spectrum is considered in isolation.  Data gathered during the pilot part of the RAVE survey has now been released \\citep{RaveDR3}, and the distances we derive from these data are contained in the release. In view of our results for the Hipparcos stars, it may be useful to correct the distances of dwarfs with $T>6000\\,$K by increasing their distances by 10\\% of their formal errors, and to correct the distances of giants by decreasing their distances by the same amount.  Our distances reveal which parts of the Galaxy the RAVE survey probes. Roughly half the stars in the RAVE catalogue are giants and half dwarfs. The giant/dwarf ratio varies strongly with Galactic latitude and to a weaker extent with Galactic longitude. The structure of the variation is accurately predicted by \\textit{Galaxia} when obscuration by dust is included, which significantly reduces the fraction of giants seen towards the Galactic Centre.  Although the spatial distribution of RAVE stars reflects the survey's selection function as well as the intrinsic stellar density of the Galaxy, it nonetheless reveals the double-exponential structure of the disc. Moreover, the distribution of stellar ages shows the expected concentration of young stars towards the plane.  The metallicity distribution evolves systematically with distance from the plane, being remarkably narrow and slightly sub-solar at $|z|<150\\pc$, to much broader and centred on $\\mh\\simeq-0.5$ more than $2\\kpc$ up. Our results support the view that observational errors have the biggest impact on the observed metallicity distribution near the plane.  This work has brought into sharp focus the crucial importance of the pipeline that extracts stellar parameters from the raw spectra. A valuable upgrade to the current pipeline would be to force the parameters assigned to a star to be consistent with models of stellar evolution. This upgrade, which is mandatory from the perspective of Bayesian inference, would eliminate the high density of stars in the bottom-right panel of Fig.~\\ref{fig:maryfig} at $\\log g\\simeq5$ and $\\Teff>6000\\K$. This study also highlights the need for a more sophisticated analysis of the errors in the parameters returned by the pipeline. For example, one would expect the error on $\\log g$ to depend on [M/H] in addition to S/N, which we have had to assume has complete control of the errors in the input parameters. Moreover, near the turnoff the errors in $\\Teff$ and $\\log g$ will be quite strongly correlated, and we have had to neglect this fact.  All improvements in the pipeline will feed through into more accurate distances.  The distances and stellar parameters described here provide a basis for extensive work on the structure, kinematics and dynamics of our Galaxy. Work on the Galaxy's kinematics is already underway, and papers in this area will appear shortly. Work on the Galaxy's dynamics requires characterisation of the intrinsic density distribution of the population(s) whose kinematics have been measured. Determining those densities involves determination of the survey's selection function. This is the next major task that must be accomplished before the RAVE survey can attain its ultimate goals."
    },
    "1107/1107.2784_arXiv.txt": {
        "abstract": "Three regions of excessive flux of cosmic rays with energies of the order of PeV are found in the experimental data of the EAS MSU array at a confidence level greater than $4\\sigma$. For two of them, there are similar regions in the experimental data of the EAS--1000 Prototype array. One of the interesting features of the regions is the absence of supernova remnants in their vicinities, traditionally considered as the main sources of Galactic cosmic rays, but the presence of isolated pulsars, some of which are able to accelerate heavy nuclei up to energies close to PeV. In our opinion, this favors the assumption that isolated pulsars are able to contribute to the flux of Galactic cosmic rays more than is usually assumed. ",
        "introduction": "The problem of the origin of cosmic rays (CRs) in the PeV energy range remains open for several decades. The model of diffusive shock acceleration at the outer front of expanding supernova remnants (SNRs) has become the most widely spread one, see, e.g.,~[1], though it has not been experimentally confirmed yet. Along with supernova remnants, other possible sources of PeV CRs have been considered. In particular, shortly after pulsars were discovered and identified with neutron stars, it was demonstrated that they can be effective accelerators of CRs~[2]. Interest to pulsars as possible sources of CRs with energies $\\gtrsim10^{14}$~eV and up to ultrahigh energies did not vanish [3--7]. Models considered included acceleration both in pulsar wind nebulae (e.g., the Crab nebula) and by isolated pulsars. It was demonstrated in particular that the Geminga pulsar is a possible candidate for being the single source of the knee at 3~PeV~[5]. There is no surprise then that one of the areas of investigations in the field are attempts to find correlation between arrival directions of CRs and coordinates of their possible astrophysical sources, including SNRs, pulsars, and some other.  In our previous works dedicated to the analysis of arrival directions of extensive air showers (EAS) registered with the EAS MSU and EAS--1000 Prototype (below, ``PRO--1000'') arrays, we have already presented a number of regions of excessive flux (REFs) of PeV cosmic rays [8--11]. It was shown that  not all of them could be matched with Galactic supernova remnants. However, a substantial part of REFs could be matched with isolated pulsars in whose vicinities nebulae are not observed. In this work, we consider three regions selected in the EAS MSU array data at a level of confidence greater than $4\\sigma$. None of these could be matched with Galactic supernova remnants, but there are isolated pulsars inside or in close vicinities of these regions, some of which are able to accelerate heavy nuclei up to energies close to PeV. For two of these regions, there are similar regions in the PRO--1000 experimental data. ",
        "conclusions": "In what follows, we consider REFs found in the EAS MSU array experimental data under the conditions \\begin{equation} S > 4,\\quad \\Pbc > 0.9545,\\quad \\Nreal > 100. \\end{equation} Three regions satisfying~(1) were found, see Fig.~1 and Table~1. Each of them consists of a number of partially overlapping cells, selected with the algorithm described above. The probability that any of the three REFs appears by chance is $\\le6\\times10^{-5}$. This estimate follows from the values of~\\PN\\ shown in Table~1, as well as from the fact that $S>4$ for all the REFs: the corresponding probability equals $\\approx0.999937$ for the significance level equal to~4. Regions~A and~B have ``partners'' in the PRO--1000 data set. These are regions that satisfy conditions $S>3$, $\\Pbc>0.9545$, $\\Nreal>100$ and partially overlap with those two regions in the EAS MSU data set.  \\begin{figure}[!ht] \\centerline{\\includegraphics[width=.7\\textwidth]{figure.eps}} \\vspace{-1pc} \\caption{Location of the regions of excessive flux in equatorial coordinates. Filled with red are the three regions found in the EAS MUS data that satisfy conditions~(1). The open red rectangle shows the boundaries of a region found with the condition $\\Pbc>0.9545$ omitted. Blue rectangles show REFs found in the PRO--1000 data set that partially overlap with the above REFs in the EAS MSU data and satisfy conditions $S>3$, $\\Pbc>0.9545$, $\\Nreal>100$ (``partners''). Filled red triangles show coordinates of the Galactic supernova remnants~\\cite{Green}. Filled blue diamonds show coordinates of the Galactic pulsars~\\cite{ATNF}. The curves at the left and right sides of the Figure show the Galactic plane, the $\\cap$-like curve shows the Supergalactic plane.} \\end{figure}  \\begin{table}[!ht] \\label{tab:MSU} \\caption{Some parameters of the REFs in the EAS MSU data set that satisfy conditions~(1) (A, B, C), and their ``partners'' in the PRO--1000 data set (A$'$, B$'$). Notation: $\\alpha$ and $\\delta$ are the ranges of values of the corresponding coordinates; \\Nreal, \\Nbg\\ are the ranges of the experimental and background flux of EAS in the cells that constitute the REF respectively, $\\max S$ and $\\min\\PN$ are the maximum value of the significance level~$S$ and the minimum value of~\\PN\\ for the cells that constitute the REF. } \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline REF & $\\alpha,\\,^\\circ$ & $\\delta,\\,^\\circ$ & \\Nreal\\ & \\Nbg\\ & $\\max S$ & $\\min\\PN$ \\\\ \\hline A & 115.5\\dots130.5 & 68\\dots72.5 & 960\\dots1498 & 841.9\\dots1350.8 & 4.233461 & 0.999962 \\\\ A$'$& 112.5\\dots136 & 71\\dots78 & 2034\\dots6406 & 1881.4\\dots6146.9 & 3.852460 & 0.999474 \\\\ \\hline B & 280\\dots287.5 & 56\\dots60 & 537\\dots1029 & 451.5\\dots904.3 & 4.478172 & 0.999958 \\\\ B$'$& 283\\dots292 & 58\\dots62.5 & 1687\\dots3473 & 1552.4\\dots3283.8 & 3.556190 & 0.999494 \\\\ \\hline C & 325\\dots329.5 & 22\\dots26.5 & 108\\dots152 & 67.34\\dots108.9 & 4.954917 & 0.999939 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The median values of~\\Ne\\ for the regions A, B, and C are close to the value obtained for the whole data set. The region that is found around region~A when omitting the condition $\\Pbc>0.9545$ lies whithin the boundaries $\\alpha=96^\\circ\\dots130.5^\\circ$ and $\\delta=67^\\circ\\dots79^\\circ$ and contains 7386 EAS. The median value of~\\Ne\\ for this REF is almost equal to the value for the whole data set. The minimum values of~\\Pbc\\ and~\\PN\\ for the cells forming this region are equal to 0.831062 and 0.999962, respectively.  As can be seen from the Figure, there are currently no known Galactic supernova remnants in the vicinity of the regions found.  Some other possible accelerators of Galactic CRs, such as X-ray binaries and OB-associations have not been found there either. There is a bright open cluster NGC 2408 at the western boundary of Region~A, but the distance from it to the Solar system is unknown, and it might be difficult to obtain a reliable estimation of its contribution to the flux of Galactic PeV CRs. On the other hand, there are isolated pulsars inside or in the close vicinities of all the regions. In view of this, the question arises as to whether these pulsars could accelerate charged particles to the energies of the order of PeV?  To the best of our knowledge, there is currently no common model of acceleration of charged particles in the vicinity of isolated pulsars. To estimate the maximum energy of a particle accelerated near the light cylinder of a pulsar, we employ an expression that can be written as follows: $E_{\\max} = 0.34 \\, Z \\, B \\, \\Omega^2$~eV, where $Z$ is the charge number of the particle, $B$ is the strength of the surface magnetic field,~G; $\\Omega$ is the angular velocity of a pulsar, rad~s$^{-1}$; and the radius of the pulsar equals $10^6$~cm~[3]. The expression was obtained under the assumption that most of the magnetic field energy in the pulsar wind zone is transformed into the kinetic energy of the particles, and the density of electron-positron pairs does not exceed $10^{-5}$ of that of ions of iron. Table~2 presents approximate values of maximum energies of protons and iron nuclei accelerated by pulsars located in the vicinities of regions~A, B, and C. It can be seen that in all three cases sufficiently heavy nuclei can be accelerated up to energies close to PeV.  \\begin{table}[!ht] \\label{tab:PSRs} \\caption{Some parameters of the pulsars located near REFs A, B, and C, and estimates of the maximum energy of protons and iron nuclei. Notation: $d$ is the distance from a pulsar to the Solar system,~kpc; $P$ is the barycentric period of rotation of a pulsar,~s; $B$ is the strength of the surface magnetic field,~G; $\\dot E$ is the rotation energy loss,~ergs/s; $\\max\\Ep$ and $\\max\\EFe$ are the maximum energies of a proton and an iron nucleus accelerated by a pulsar,~eV.} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|c|c|} \\hline REF& PSR Name & $d,$ & $P,$  & Age,              & $B,$     & $\\dot E,$ & $\\max\\Ep,$ & $\\max\\EFe,$ \\\\ &            & kpc  &  s    & year              &  G       & ergs/s    &  eV        & eV \\\\ \\hline A& J0814+7429 & 0.33--0.43 & 1.292 & $1.22\\cdot10^{8}$ & $4.72\\cdot10^{11}$ & $3.08\\cdot10^{30}$ & $4\\cdot10^{12}$ & $1\\cdot10^{14}$  \\\\ & J0700+6418 & 0.48 & 0.196 & $4.52\\cdot10^{9}$ & $1.17\\cdot10^{10}$ & $3.61\\cdot10^{30}$ & $4\\cdot10^{12}$ & $1\\cdot10^{14}$  \\\\ & J0653+8051 & 3.37 & 1.214 & $5.07\\cdot10^{6}$ & $2.17\\cdot10^{12}$ & $8.37\\cdot10^{31}$ & $2\\cdot10^{13}$ & $5\\cdot10^{14}$  \\\\ & J0849+8028 & 3.38 & 1.602 & $5.69\\cdot10^{7}$ & $8.56\\cdot10^{11}$ & $4.28\\cdot10^{30}$ & $4\\cdot10^{12}$ & $1\\cdot10^{14}$  \\\\ \\hline B& J1836+5925 &0.17--0.75& 0.173 & $1.80\\cdot10^{6}$ & $5.20\\cdot10^{11}$ & $1.16\\cdot10^{34}$ & $2\\cdot10^{14}$ & $6\\cdot10^{15}$  \\\\ & J1840+5640 & 1.70 & 1.653 & $1.75\\cdot10^{7}$ & $1.59\\cdot10^{12}$ & $1.31\\cdot10^{31}$ & $8\\cdot10^{12}$ & $2\\cdot10^{14}$  \\\\ \\hline C& J2155+2813 & 5.06 & 1.609 & $2.78\\cdot10^{7}$ & $1.23\\cdot10^{12}$ & $8.68\\cdot10^{30}$ & $6\\cdot10^{12}$ & $2\\cdot10^{14}$  \\\\ & J2156+2618 & 4.71 & 0.498 & $5.56\\cdot10^{8}$ & $8.51\\cdot10^{10}$ & $4.53\\cdot10^{30}$ & $5\\cdot10^{12}$ & $1\\cdot10^{14}$  \\\\ & J2151+2315 & 1.42 & 0.594 & $1.33\\cdot10^{7}$ & $6.56\\cdot10^{11}$ & $1.34\\cdot10^{32}$ & $2\\cdot10^{13}$ & $6\\cdot10^{14}$  \\\\ & J2139+2242 & 4.71 & 1.084 & $1.21\\cdot10^{7}$ & $1.26\\cdot10^{12}$ & $4.41\\cdot10^{31}$ & $1\\cdot10^{13}$ & $4\\cdot10^{14}$  \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The distances to the pulsars from Table~2 exceed considerably the free path of PeV neutrons, and the amount of electromagnetic showers is insufficient for formation of REFs. In view of this, the question arises of how to keep the arrival direction of (charged) CRs propagating through the Galactic magnetic field.  We are currently not ready to suggest a satisfactory model for answering this question, but we note at least two models were suggested to explain the REFs of CRs with energies of the order of 10~TeV registered with the Milagro gamma-ray observatory~[18]. One of the models is based on CR acceleration in the vicinity of the Geminga pulsar [19], the properties of which are similar to those of the pulsar J1836+5925 located at the boundary of the B region~[20]. The other one uses a specific magnetic field configuration able to focus a flux of charged particle~[21].   To conclude, we have presented three regions of excessive flux of PeV CRs selected from the EAS MSU experimental data at a significance level $>4\\sigma$ and satisfying several additional conditions. In the vicinity of every region, there are isolated Galactic pulsars that are able to accelerate sufficiently heavy nuclei to the required energies. In our opinion, this does not give enough ground to claim that the above REFs are formed due to contribution from these pulsars but can witness in favor of the assumption that (1) some isolated pulsars contribute to the flux of PeV CRs more than is usually assumed and (2) there are efficient mechanisms of focusing CR in the interstellar space.  \\bigskip  Only free, open source software was used for the investigation. In particular, all calculations were performed with GNU Octave~[22] running in Linux. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France (\\texttt{http://simbad.u strasbg.fr/simbad}).  The research was partially supported by the Russian Foundation for Fundamental Research grant No.~08-02-00540"
    },
    "1107/1107.1672_arXiv.txt": {
        "abstract": "{The dip model  assumes that the ultra-high energy cosmic rays (UHECRs) above 10$^{18}$ eV consist  exclusively of protons and  is consistent with the spectrum and composition measure by HiRes.  Here we present the range of cosmogenic neutrino fluxes  in the dip-model which are compatible with a recent determination of the extragalactic very high energy (VHE)  gamma-ray diffuse background derived from 2.5 years of Fermi/LAT data. We show that  the largest fluxes predicted in the dip model would be detectable by IceCube in about 10 years of observation and are  within the reach of a few years of observation with the ARA project. In the incomplete UHECR model in which protons are assumed to dominate only above 10$^{19}$ eV, the cosmogenic neutrino fluxes could be a factor of 2 or 3 larger. Any fraction of  heavier nuclei in the UHECR at these energies would reduce the maximum cosmogenic neutrino fluxes.  We also consider here special evolution models in which the UHECR sources are assumed to have the same evolution of either the star formation rate (SFR),  or the gamma-ray burst (GRB) rate, or the active galactic nuclei (AGN) rate in the Universe and found that the last two are disfavored (and in the dip model rejected) by the new VHE gamma-ray background. }   \\begin{document} ",
        "introduction": "The origin and composition of the ultra-high energy cosmic rays (UHECRs) remain  open questions in spite of the experimental advances of recent years due to the Auger  and HiRes collaborations.  The Auger collaboration~\\cite{Auger} has found evidence in its measurements of air shower profiles, both in the depth of shower maximum  and its fluctuation, of an increasing dominance of heavy nuclei primaries at high energies~\\cite{xmax},  starting with light nuclei at $3 \\times 10^{18}$~eV and approaching iron at $4 \\times 10^{19}$.   The interpretation of the air shower profiles depends, however, on the hadronic interaction  models used, which have uncertainties at large energies. Auger has also found  a correlation of events above $ 5.5 \\times 10^{19}$ eV with the distribution of nearby extragalactic astrophysical sources, including an excess around the direction of Centauros A, the nearest AGN. The correlation of UHECR arrival directions with  the distribution of relatively close AGN within an angular distance of 3$^o$ has weakened with time~\\cite{correlation}. This correlation is compatible with proton primaries, which would contradict the  heavy composition extracted from the Auger shower profiles. However, it has been argued that the  excess of events  in the direction of Centaurus A, the most significant contribution to the anisotropy in the Auger data,  could be explained by heavy nuclei primaries emitted from the Virgo cluster~\\cite{Semikoz-Virgo}.  No tendency towards a heavier composition at high energies is found in the HiRes~\\cite{HiRes} data, which are compatible with the UHECR above 10$^{19}$~eV consisting primarily of   protons~\\cite{Abbasi:2009nf}. HiRes data seem to indicate a change in composition from heavy to light at energies close to 5$\\times 10^{17}$ eV~\\cite{chem_HiRes}. However, HiRes does not find  a correlation of the  arrival directions  of UHECR above $4.0 \\times 10^{19}$ eV or $ 5.7 \\times 10^{19}$ eV with astrophysical sources at a distance smaller than 10$^o$~\\cite{Abbasi:2010xt}.  The spectrum and composition measured by HiRes are, thus, consistent with assuming that all UHECR above $10^{18}$ eV are due to  extragalactic protons. This is what is assumed in the ``dip model\"~\\cite{dip-model}. In the past we called this  the ``minimal UHECR model\"~\\cite{minimal-GKS}, because it allows to  fit the UHECR spectrum above 10$^{18}$ eV with   a minimum number of parameters. It is  usually called the ``dip-model\", because  the ankle feature in the spectrum at $\\sim 4 \\times 10^{18}$ eV is interpreted as an absorption ``dip\" in the spectrum due to energy losses by electron-pair production of the protons interacting with the cosmic microwave background (CMB).  Both HiRes~\\cite{Abbasi:2007sv} and  Auger~\\cite{PAO-spectrum}   have  observed a suppression in the spectrum above $4 \\times 10^{19}$~eV. For HiRes this is  consistent with being due to the photo-pion production of  protons interacting with the CMB, i.e. the GZK feature~\\cite{gzk}.  The GZK process produces pions, the decay of which produces  both ``cosmogenic neutrinos\"~\\cite{bere} (from $\\pi^{\\pm}$)  and   photons (from  $\\pi^0$), which we  call ``GZK photons\". Although they share the same production mechanism, ultra-high energy (UHE) GZK photons and cosmogenic neutrinos are affected very differently by intervening backgrounds. The flux of UHE GZK photons is affected by the poorly known radio background and extragalactic magnetic fields which do not affect  neutrinos. These photons only reach us from less than 100 Mpc away, i.e. at redshifts $z<0.02$ (see e.g.~\\cite{GKS}). Auger has placed stringent limits on the fraction of UHECRs that are photons~\\cite{augersdphotonfractionlimit}. Cosmogenic neutrinos do not interact during propagation and thus reach us from the whole production volume. Thus the flux of cosmogenic neutrinos arriving to Earth depends strongly on the evolution of the sources, which affect mostly the density of sources far away.  The cosmogenic neutrinos have been extensively studied theoretically  since 1969~\\cite{bere} onwards  (see e.g.~\\cite{reviewGZKneutrinos,Semikoz:2003wv} and references therein), and constitute one of the main high energy signals expected in neutrino telescopes, such as IceCube~\\cite{ICECUBE},  ANITA~\\cite{ANITA},  the future KM3NeT~\\cite{KM3NeT} and ARA~\\cite{ARA} or space based observatories such as JEM-EUSO~\\cite{JEM-EUSO}. They could also be observed by  the Pierre Auger  Observatory~\\cite{Auger} and the Telescope Array~\\cite{TA}.  We want here to see if  the highest cosmogenic neutrino fluxes compatible with all present data are within the reach of these experiments. The expected flux of cosmogenic neutrinos is higher if the UHECR consist of protons.  Heavy  nuclei  would  interact with radiation backgrounds and  break up into lighter nuclei and nucleons, only a fraction of which would be above the energy threshold for pion-production. Therefore  fewer UHE neutrinos (and  GZK photons)  would be produced. Thus we assume proton primaries and, to be consistent with this assumption, use the UHECR spectrum  measured by HiRes~\\cite{HiRes-spectrum, Abbasi:2007sv}, which is also higher than the Auger spectrum (thus it also leads to  larger secondary neutrino fluxes because  of this reason). Thus we assume the  dip model~\\cite{dip-model}, which is self consistent model of UHECR. Because we are interested in producing upper bounds to neutrinos fluxes, we also consider the case in which protons only fit the HiRes spectrum above 10$^{19}$ eV, although this is  an incomplete model which does not explain the UHECR data below this energy. This model was also studied in recent evaluations of cosmogenic neutrino fluxes, Refs.~\\cite{Berezinsky:2010xa}  and \\cite{Ahlers:2010fw}, where it was found to produce the highest maximum fluxes.  Following  Ref.~\\cite{Berezinsky:2010xa}  we call it the ``ankle model\".  In this paper we  present  the  range of cosmogenic neutrinos in the dip and ankle models obtained by varying all relevant parameters  defining the primary astrophysical extragalactic nucleon fluxes  with which the HiRes data are well fitted above  either 10$^{18}$ eV or 10$^{19}$ eV, respectively. Our code  uses the  kinetic equation approach, which  provides fast results for each of the set of parameters we use to predict each flux. This feature allows us to scan the  parameter space  to look for the range of predicted  fluxes. We already studied the expected range of GZK photons in the dip model~\\cite{minimal-GKS}. Here we  study the expected range of cosmogenic neutrinos using a similar statistical method and taking into account new data, in particular the extragalactic  flux of photons below 1 TeV extracted from measurements of the Fermi Space Telescope ~\\cite{Abdo:2010nz, Neronov:2011kg}.  The  secondary electrons, positrons and photons quickly cascade on the CMB, infrared and visible backgrounds  to lower  energies and generate a $\\gamma$-ray background in the GeV-TeV energy range, at which  the Universe becomes transparent to photons~\\cite{wdowczyk}. We explore in particular the constraint imposed on the neutrino production  models  by a recent evaluation~\\cite{Neronov:2011kg}  of the diffuse extragalactic very high energy (VHE) gamma-ray background using 2.5 years of Fermi/LAT data. ",
        "conclusions": "Cosmogenic neutrinos constitute one of the main high energy signals expected in neutrino telescopes, such as IceCube~\\cite{ICECUBE},  ANITA~\\cite{ANITA},  the future KM3NeT~\\cite{KM3NeT} and ARA~\\cite{ARA} or space based observatories such as JEM-EUSO~\\cite{JEM-EUSO}.  They could also be observed by Pierre Auger Observatory~\\cite{Auger} and the Telescope Array~\\cite{TA}. Here we present the largest neutrino fluxes compatible with cosmic rays and gamma-rays observations.  Proton primaries produce larger maximum cosmogenic neutrino fluxes than heavier nuclei. Thus we consider the  dip model, which assumes that all UHECR primaries above 10$^{18}$ eV are protons  and is consistent with both the spectrum and composition data of HiRes. As it  has been  already pointed out~\\cite{Berezinsky:2010xa, Ahlers:2010fw} fitting the UHECR spectrum  with protons only above  higher energies, such as 10$^{19}$ eV, leads to higher neutrino fluxes. This, however, is an incomplete model of UHECR unless a consistent explanation of the dominant component below this energy is presented. With this caveat, we  also consider this model, which following Ref.~\\cite{Berezinsky:2010xa} we call the ankle model, to obtain the highest allowed cosmogenic neutrino fluxes.  We use a program which allows for fast analytic calculations of the secondary proton, photon and neutrino spectrum to find at each energy the expected range of cosmogenic neutrino fluxes. We require the produced UHECR spectrum to fit the HiRes spectrum at energies above a certain energy $E_{\\rm fit}$, which is 10$^{18}$ eV in the dip model and 10$^{19}$ eV in the ankle model. Below this energy, the HiRes spectrum is taken  as an upper bound. We also impose as an upper bound to the  predicted gamma-ray fluxes  the new evaluation of the diffuse extragalactic VHE $\\gamma$-ray background of Ref.~\\cite{Neronov:2011kg} and compare the results with those derived assuming instead the First-Year Fermi~\\cite{Abbasi:2007sv} background, which is higher than the former by about a factor of 2 at energies above 10 GeV.  The resulting neutrino flux ranges are presented in Figs.~\\ref{F6} to \\ref{F9} together with experimental bounds of IceCube-40, Auger and ANITA and sensitivity regions  of  IceCube in 3 years, ARA in 3 years and JEM-EUSO.  In the left panels of Figs.~\\ref{F6}, \\ref{F7} and \\ref{F8}  we show the range of cosmogenic neutrino fluxes as function of the energy expected in the dip-model of UHECR allowing for a shift of the HiRes energy scale by a factor between $(1.30)^{-1}$ and 1.30, to take into account possible systematic errors. Also in these panels the results obtained by taking as upper bounds both VHE $\\gamma$-ray backgrounds are compared.  The  comparison of the ranges of neutrino fluxes obtained in the dip model with and without shifting the nominal HiRes energy are presented in the right  panels of Figs.~\\ref{F6}, \\ref{F7} and \\ref{F8}  for the particular case yielding the largest fluxes with the new VHE $\\gamma$-ray spectrum of Ref.~\\cite{Neronov:2011kg} in the corresponding left panel of the same figure. The ranges obtained varying the energy are larger than with the energy fixed, but not by much, because the HiRes spectrum without any energy shift is very well fitted by the dip model, so large changes in the HiRes spectrum do not provide good fits.  In  Figs.~\\ref{F6} and \\ref{F7}, the evolution model in Eq.~\\ref{evolution-function} is used.  In these figures, the highest cosmogenic neutrino fluxes allowed by the lower estimate of the VHE $\\gamma$-ray background~\\cite{Neronov:2011kg} are about a factor of 2 lower that those obtained by imposing the First-Year Fermi background.  Figs.~\\ref{F6}, \\ref{F7} and \\ref{F8} show that  detecting the highest cosmogenic neutrino fluxes predicted by the dip model would require about 10 years of observation with IceCube (or a similar KM3NeT) or the construction of the new ARA detector.  The upper and middle panels of Fig.~\\ref{F9} are similar to the left panels of Figs.~\\ref{F6} and \\ref{F7} but apply to the  ankle model. Here the maximum cosmogenic neutrino fluxes are a factor of 2 to 3 higher than in the dip model, what brings them into the range of detectability of IceCube in 3 years. Our conclusions are consistent with those of Ref.~\\cite{Ahlers:2010fw}  in that the ankle model produces maximal neutrino fluxes which could be detected with a few years of observation of IceCube, and are not consistent with those of Ref.~\\cite{Berezinsky:2010xa} in this regard.  Refs.~\\cite{Berezinsky:2010xa} and \\cite{Ahlers:2010fw} use the evolution model in Eq.~\\ref{evolution-function}. We also consider here the special evolution functions  in Eqs.~~\\ref{SFR}, \\ref{GRB} and \\ref{AGN}, in which the UHECR sources are assumed to have the same evolution of either the star formation rate (SFR),  or the gamma-ray burst (GRB) rate, or the active galactic nuclei (AGN) rate in the Universe.  To our knowledge the limits imposed by the diffuse extragalactic VHE $\\gamma$-ray background extracted from Fermi-LAT data on these evolution functions had not been studied before.  Fig.~\\ref{F8} shows that the new  VHE $\\gamma$-ray background~\\cite{Neronov:2011kg} rejects the dip models assumed to have the GRB source evolution function in Eq.~\\ref{GRB}. The dip models with the AGN source evolution function in Eq.~\\ref{AGN} are rejected by any of the two VHE $\\gamma$-ray backgrounds, so models with this evolution do not appear in the figure. Only dip models with the SFR evolutions function in Eq.~\\ref{SFR} is allowed by both observed  VHE $\\gamma$-ray backgrounds assumed, and both ranges (only the maximum fluxes change) are very similar. The lower panel of Fig.~\\ref{F9}  shows that in the ankle model,  both the GRB and AGN  source evolution functions are allowed, even with the lower $\\gamma$-ray background we used (in the case of the AGN evolution barely so).  \\vspace{0.3cm} {\\bf Acknowledgments} The numerical calculations in this work were performed with the Computational Cluster of the Theoretical Division of INR RAS. G.G was supported in part by the US DOE grant DE-FG03-91ER40662 Task C. O.K. was supported in part by the RFBR grant 10-02-01406-a. G.G. wants to thank the APC, Paris, for its hospitality  during part of the period in which this paper was written."
    },
    "1107/1107.3394_arXiv.txt": {
        "abstract": "We study how the presence of a background magnetic field, of intensity compatible with current observation constraints, affects the linear evolution of cosmological density perturbations at scales below the Hubble radius. The magnetic field provides an additional pressure that can prevent the growth of a given perturbation; however, the magnetic pressure is confined only to the plane orthogonal the field. As a result, the ``Jeans length'' of the system not only depends on the wavelength of the fluctuation but also on its direction, and the perturbative evolution is anisotropic. We derive this result analytically and back it up with direct numerical integration of the relevant ideal magnetohydrodynamics equations during the matter-dominated era. Before recombination, the kinetic pressure dominates and the perturbations evolve in the standard way, whereas after that time magnetic pressure dominates and we observe the anisotropic evolution. We quantify this effect by estimating the eccentricity $\\epsilon$ of a Gaussian perturbation in the coordinate space that was spherically symmetric at recombination. For a perturbations at the sub-galactic scale, we find that $\\epsilon = 0.7$ at $z=10$ taking the background magnetic field of order $10^{-9}$ gauss. ",
        "introduction": "Our theoretical knowledge of the Universe is based on the Standard Cosmological Model, that provides a convenient framework to satisfactorily explain the majority of cosmological observations, like the anisotropy pattern of the cosmic microwave background (CMB) \\cite{Komatsu:2010fb,Larson:2010gs,Dunkley:2010ge}, the large-scale structure of the Universe \\cite{Tegmark:2004,Cole:2005sx,Tegmark:2006az}, the Hubble diagram of distant type Ia supernovae \\cite{Riess:1998cb,Perlmutter:1998np,Frieman:2008sn} and the abundances of light elements \\cite{Iocco:2008va}.  The Standard Cosmological Model relies on the assumption that the Universe, at least at large scales, is highly homogeneous and isotropic, and its geometry is thus described by the Robertson-Walker metric. In fact, the distribution of luminous red galaxies shows that the present Universe is homogeneous on scales greater than $\\sim 100$ Mpc \\cite{Hogg:2004vw}, while the isotropy of the CMB itself (which has a black-body distribution at $T=2.73\\,$K with temperature fluctuations of order $10^{-5}$ or less) is an indication of the isotropy of the Universe as a whole and a strong evidence for homogeneity at the time of hydrogen recombination (nearly 400.000 years after the Big Bang, corresponding to a cosmological redshift $z_\\mathrm{rec}=1100$). On the other hand, below the ``homogeneity scale'' of 100 Mpc, the distribution of matter is definitely inhomogeneous. Such a dichotomy between the smoothness in the matter-energy distribution at $z=\\zrec$ and the clumpiness of the recent Universe (for $z\\lesssim1$) below a certain scale is explained by the mechanism of gravitational instability: the structures we observe today have been formed through the growth of tiny density perturbation seeds that, accordingly to the currently accepted model, were created in the early Universe during a phase of inflationary expansion.  The presence of a large scale (\\ie coherent over a Hubble length), strong magnetic field is forbidden by the observed isotropy, as it would naturally single out a preferred spatial direction. However, a background and uniform magnetic field could be present at cosmological scales provided that its intensity is small enough. In particular, upper limits on the present field intensity of order $\\sim 10^{-9}$ G have been derived from observations of the CMB temperature anisotropies \\cite{BFS97,Paoletti:2010rx} and of its temperature-polarization correlation \\cite{Scannapieco:1997mt,Komatsu:2010fb}\\footnote{It has been argued (see, e.g., \\cite{Adamek:2011pr}) that the limits obtained from the  CMB temperature data can be significantly relaxed in the presence of free-streaming neutrinos. However this does not affect the limits from the polarization.}. Smaller scale fields, on the other hand, could be as strong as $10^{-6}$ G.  The effects of large-scale magnetic fields on the evolution of cosmological structures have been studied extensively in the literature (for a complete review, see Ref. \\cite{BMT07} and references therein), where both Newtonian and general relativistic treatments, the latter often using covariant and gauge-invariant techniques, are present. It is known that, among others, magnetic fields slow down (and possibly prevent) the growth of perturbations and can produce vorticities and shape distortions in the density field \\cite{VTP05,Tsagas:1999ft}. In this Letter, our goal is to revisit the issue of the existence of a ``magnetic Jeans length'' and of its dependence on the direction along which the perturbation propagates, other than to give a realistic estimate of its value in a matter-dominated Universe. The presence of a magnetic Jeans length has been discussed in several papers \\cite{BMT07,VTP05,Tsagas:1999ft,RR71, Kim:1994zh} both in a Newtonian and in general relativistic framework (for an analysis of the standard Jeans mechanism in the presence of dissipative effect see also \\cite{CQG,IJMPD,MPLA}), but some of the analyses failed to recognize its angular dependence. Here, we present a neat derivation of the relevant instability scales bases on the equations of magnetohydrodynamics (MHD) on an expanding Universe. We discuss a simple generalization of the magnetic Jeans length suited for two-fluid systems, and clarify some misunderstandings that are present in the literature. We also show the results of the numerical integration of the coupled MHD and Poisson equations. Finally, we numerically study the distortion introduced in the density field by the anisotropy in the critical length.  The Letter is organized as follows. In Sec. 2, we characterize the Universe as a plasma. In Sec. 3, we introduce the basic equations, and in Sec. 4 we carry out the linearization procedure. We derive the existence of the magnetic Jeans length from analytical considerations in Sec. 5, while in Sec. 6 we show some numerical results. Finally, we draw our conclusions in Sec. 7. ",
        "conclusions": "We have studied the effect of a background magnetic field on the linear evolution of cosmological density perturbations at scales well below the Hubble length, where a Newtonian treatment can be used, focusing on the matter-dominated era. The conditions that allow for the growth of small density perturbations have been clearly stated. In particular, we have found that in the plane orthogonal to the ambient magnetic field, a new critical length appears, related to the presence of the magnetic pressure, while everywhere else outside that plane the stability is dictated by the standard Jeans criterion. This is also confirmed through a direct numerical integration of the relevant MHD equations during the matter-dominated era, and this effect is shown to be possibly important after recombination, when the magnetic pressure of baryons is much larger than their the kinetic pressure. Finally, it has been shown how the dependence of the critical scale on the angle between the perturbation wavevector and the magnetic field could lead to a sizable anisotropy in the perturbations at sub-galactic scales at the onset of non-linearity.  Our analysis has relied on some approximations: in particular, we have ignored the gravitational effects of dark matter perturbations. We have argued that this approximation limit the validity of our treatment to some time after recombination. We defer a more detailed and fully general relativistic analysis, also taking into account the different fluid components, to a future work.  {\\small \\textbf{Acknowledgment}: N.C. gratefully acknowledges the \\emph{CPT - Universit\\'e de la Mediterran\\'ee Aix-Marseille 2} and the financial support from ``Sapienza'' University of Roma. M.L. acknowledges  support from a joint \\emph{Accademia dei Lincei / Royal Society} fellowship for Astronomy. The work of M.L. has been supported by \\emph{MIUR - Ministero dell'Istruzione, dell'Universit\\`a e della Ricerca} through the PRIN grants ``Matter-antimatter asymmetry, Dark Matter and Dark Energy in the LHC era'' (contract number PRIN 2008NR3EBK-005) and \"Galactic and extragalactic polarized microwave emission\" (contract number PRIN 2009XZ54H2-002).}    \\newpage \\appendix"
    },
    "1107/1107.3993_arXiv.txt": {
        "abstract": "We have analyzed the new {\\it Rossi X-ray Timing Explorer} Proportional Counter Array data of the atoll neutron star (NS) low-mass X-ray binary (LMXB) system XB 1254--690. The colour-colour diagram shows that the source was in the high-intensity banana state. We have found two low-frequency candidate peaks with single trial significances of $\\approx 2.65 \\times 10^{-8}$ and $\\approx 7.39 \\times 10^{-8}$ in the power spectra. After taking into account the number of trials, the joint probability of appearance of these two peaks in the data set only by chance is $\\sim 4.5\\times10^{-4}$, and hence a low-frequency QPO can be considered to be detected with a significance of $\\sim 4.5\\times10^{-4}$, or, $\\sim 3.5\\sigma$ for the first time from this source. We have also done the first systematic X-ray spectral study of XB 1254--690, and found that, while one-component models are inadequate, three-component models are not required by the data. We have concluded that a combined broken-powerlaw and Comptonization model best describes the source continuum spectrum among 19 two-component models. The plasma temperature ($\\sim 3$ keV) and the optical depth ($\\sim 7$) of the Comptonization component are consistent with the previously reported values for other sources. However, the use of a broken-powerlaw component to describe NS LMXB spectra has recently been started, and we have used this component for XB 1254--690 for the first time. We have attempted to determine the relative energy budgets of the accretion disc and the boundary layer using the best-fit spectral model, and concluded that a reliable estimation of these budgets requires correlations among time variations of spectral properties in different wavelengths. ",
        "introduction": "The persistent neutron star low mass X--ray binary (LMXB) system XB 1254--690 exhibits energy dependent intensity dips with the binary orbital period ($\\sim 3.88$ hr), thermonuclear X-ray bursts, and flares \\citep{Griffithsetal1978, Courvoisieretal1986, Masonetal1980, Smaleetal2002}. This is a so-called ``atoll\" source \\citep{Bhattacharyya2007}. An atoll source typically traces a `C'-like curve in its colour-colour diagram or CCD \\citep{vanderKlis2006}. The high-intensity lower part of this curve is known as the banana state \\citep{vanderKlis2006}. XB 1254--690 has always been found in the banana state, and has never been observed in the spectrally hard low-intensity island state \\citep{Bhattacharyya2007}. Note that the position of a source on the CCD is usually correlated with the timing features observed from the source \\citep{vanderKlis2006}. Therefore, CCDs are useful to determine the source state, and hence to understand the related physics.  Fast brightness oscillations, i.e., burst oscillations, have been observed from thermonuclear bursts from several neutron star LMXBs (\\citet{StrohmayerBildsten2006}; \\citet{Bhattacharyya2010} and references therein). This timing feature provides one of the two ways to measure neutron star spin frequencies for LMXB systems \\citep{Chakrabartyetal2003, Bhattacharyya2010}. \\citet{Bhattacharyya2007} reported a suggestive evidence of 95 Hz oscillations from a thermonuclear burst from XB 1254--690. This indicates that the neutron star in this source spins with a frequency of 95 Hz, although further confirmation is required to establish this.  Various types of quasi-periodic oscillations (QPOs) are observed from neutron star LMXBs \\citep{vanderKlis2006}. Their frequencies are roughly between millihertz (mHz) and kilohertz (kHz). Properties of one type of QPOs are sometimes correlated with those of another, and typically these properties are correlated with the spectral state inferred from CCD. Therefore, the QPOs can be very useful to understand the accretion process around a neutron star. In addition, kHz QPOs can be useful to probe the strong gravity regime, and to measure the neutron star parameters \\citep{vanderKlis2006, Bhattacharyya2010}. However, any detection of a QPO from XB 1254--690 was never reported before this paper.  The continuum spectrum of a neutron star LMXB is typically well described with various models, giving different physical interpretations \\citep{Linetal2007}. For example, two classical spectral models, {\\it Eastern} and {\\it Western}, have been proposed for the high-intensity states \\citep{Mitsudaetal1989, Whiteetal1988}. The {\\it Eastern} model consists of a multicolour blackbody from the disc and a Comptonized component from the boundary layer (the contact layer between the neutron star and the disc inner edge), while the {\\it Western} model has a Comptonized emission from the disc and a single temperature blackbody from the boundary layer. Although further detailed studies have been done with more sources (e.g., \\citet{ChurchBalucinskaChurch2001, ChristianSwank1997, MaccaroneCoppi2003, MaitraBailyn2004, Gilfanovetal2003, Oliveetal2003, Wijnands2001, Barret2001}), a general consensus on the appropriate X-ray spectral model has not been achieved. It is therefore essential to do systematic spectral studies of neutron star LMXBs using all the reasonable physical models. Even before going into the details, the first aim of such a study will be to identify which spectral component originates from the boundary layer, and which one comes from the disc, and hence to find out the relative energy budgets of these two X-ray emitting components. These relative budgets not only will be useful to understand the accretion process and components, but also will be helpful to constrain the neutron star parameters and equation of state (EoS) models \\citep{Bhattacharyyaetal2000, Bhattacharyya2002, Bhattacharyya2010}. Recently, \\citet{Linetal2007} have done such a systematic study for two neutron star LMXBs (Aql X--1 and 4U 1608--52). However, such a study has never been done for XB 1254--690.  In this paper, we report the first evidence of a QPO from XB 1254--690. Moreover, we give the results and interpretations of the first systematic X-ray spectral study of this source. In \\S~\\ref{DataAnalysisandResults}, we mention the spectral and timing analysis methods, and report the results. In \\S~\\ref{Discussion}, we discuss the interpretation and importance of the observational results. ",
        "conclusions": "\\label{Discussion}  In this paper, we have studied the timing and spectral properties of the neutron star LMXB XB 1254--690. This atoll source was in the banana state during all the {\\it RXTE} observations. We have not detected burst oscillations from the two bursts in the analyzed data set. This is not surprising, as only a fraction of bursts from burst oscillation sources show oscillations \\citep{Galloway2008}. However, this somewhat reduces the significance of the plausible burst oscillation feature reported in \\citet{Bhattacharyya2007} by increasing the number of trials. We report the first evidence of QPOs from this source, with frequencies consistent with the $\\sim 30-80$~Hz QPOs typically observed from atoll sources in intermediate to high states \\citep{vanderKlis2006}.  We have also done the first systematic study of X-ray continuum spectrum from XB 1254--690. We have fitted an LB spectrum and a UB spectrum with 19 models (see \\S~\\ref{SpectralAnalysis} and Table 1). In both cases, a broken-powerlaw plus a Comptonization model ({\\ttfamily bknpower + comptt}) is favoured, and the plasma temperature and the optical depth of the Comptonization component are consistent with the previously reported values for other sources \\citep{Linetal2007}. A broken-powerlaw component was never used to fit the XB 1254--690 continuum spectrum, although it has been recently successfully used for the spectral fitting of Aql X--1 and 4U 1608--52. This component can represent Comptonization, either under complex conditions or in combination with another radiation process (e.g., synchrotron; see \\S~\\ref{SpectralAnalysis}). These imply that the blackbody radiations from the disc and the boundary layer are entirely processed before reaching the observer. Note that the classical models {\\it Eastern} and {\\it Western} give a much worse fit than the {\\ttfamily bknpower + comptt} model. For example, $\\chi_\\nu^2$ (dof) values for {\\it Eastern}, {\\it Western} and {\\ttfamily bknpower + comptt} models are 1.96 (30), 1.50 (30) and 0.54 (28) respectively for the LB spectrum. Besides, the {\\ttfamily cutoffpl + bbody} and {\\ttfamily diskbb + powerlaw} models, which were previously used for XB 1254--690 \\citep{Smaleetal2002, Iariaetal2007}, give much worse fits ($\\chi_\\nu^2$ (dof): 1.65 (31) and 1.95 (32) respectively for LB). However, although a broken-powerlaw component gives the best fit, note that it is only a phenomenological model, and should be replaced with a physical model in the future in order to clearly understand the meaning of the best-fit parameter values. A physical model will also be necessary to interpret the high $\\Gamma_2$ of the UB broken-powerlaw (Table \\ref{UBmodel2}), which indicates a quick cut-off after the first powerlaw component.  With the best-fit model in hand, we have attempted to determine the relative energy budgets of the disc and the boundary layer. For the LB spectrum, the flux from the {\\ttfamily bknpower} component is about 2.1 times larger than the {\\ttfamily comptt} component (Table 2). Therefore, assuming one component entirely originates from the disc and the other component fully comes out of the boundary layer, we can make the following conclusions using the Fig. 5 of \\citet{Bhattacharyyaetal2000}. (1) The spin frequency depends on the chosen EoS model of the neutron star, with larger frequencies for softer EoS models. Therefore, an independent measurement of the stellar spin frequency can be useful to constrain the EoS models. (2) If the {\\ttfamily bknpower} component originates from the disc, then the neutron star is spinning with a frequency close to the break-up frequency (irrespective of the EoS model). However, these simple conclusions may not be reliable for one or more of the following reasons. (1) Each of the spectral components may originate partially from the disc, and partially from the boundary layer. (2) Disc emission and/or boundary layer emission may be partially obscured. (3) Boundary layer emission (and also the inner disc emission) may be partially reprocessed by the disc. Besides, the {\\ttfamily bknpower} flux is smaller than the {\\ttfamily comptt} flux for the lower quality UB spectrum, although the spectral shapes and parameters of these two components are very similar to those of the LB spectrum (see Tables 3 and 4; Figs.~\\ref{specLB} and \\ref{specUB}). This difference in the normalization ratio of the two components may be due to one or more of the above three reasons. One main problem in identifying and understanding the emissions from various source components is the lack of simultaneous broadband (optical, UV and X-ray) data. The optical and UV emissions are expected to originate by the reprocessing of the X-ray photons at the outer disc, and hence the correlations among time variations of spectral properties (e.g., \\citet{Balucinskaetal2004}), especially in different wavelengths will significantly contribute to the understanding of the neutron star LMXB components. The proposed {\\it Astrosat} satellite will be useful for such a study in the near future."
    },
    "1107/1107.1444_arXiv.txt": {
        "abstract": "Taking into account a range of parameters determined from the evolutionary models and available observational data, the detailed non-LTE spectrum for the primary star and  the irradiated donor star in the AM CVn system SDSS J0926+3624 are constructed based on the TLUSTY stellar atmosphere code. The combined spectrum of the primary and the donor stars along with a multi-color blackbody spectrum of the accretion disk that reproduces a detailed numerical model is compared to the SDSS optical spectrum of the system. The photometric flux of the primary star inferred from eclipse observations is compared with the synthetic spectrum. The model fit of the two independent observations provides an upper limit on the distance of the system for different effective temperatures of the primary. In addition, an upper limit on the combined flux of the disk and the donor in the infrared region wherein the contribution of the primary is negligible is also determined.  It is shown that the spectrum of a sufficiently cool donor can exhibit emission lines due to irradiation from a hot primary and the emission features should be detectable in the infrared even though the contribution of the flux from the disk dominates. Thus, it is pointed out that infrared observations of the system would provide important information on the thermal state of the donor as well as useful insight on the thermal properties of the primary star and the accretion disk. ",
        "introduction": "The only known eclipsing AM Canum Venaticorum object to date, SDSS J0926+3624 \\citep{and05}, belongs to the class of ultra compact binary systems characterized by a pair of interacting white dwarfs with the shortest orbital period among any other binary subclass.  Such a system is driven by the loss of orbital angular momentum by the emission of gravitational waves and is among a class of systems thought to constitute the largest fraction of interacting ultra compact binaries in the Galaxy [see \\cite{nel01,nel04}]. An understanding of their properties is important for their detectability not only as individual candidate sources, but also as background sources of low frequency gravitational waves [e.g.,\\cite{nel04,sto05}].  A detailed understanding of the formation of SDSS J0926+3624 (hereafter J0926+3624) also allows one to place important constraints on the efficiency of the common envelope phase of binary evolution in forming short period binary systems containing double white dwarf components [e.g., \\cite{nel05}; for a review of the common envelope phase of evolution for the general binary population, see \\cite{ibe93,taam00}]. For AM CVn systems, in general, their formation may involve a white dwarf or helium star channel \\citep{nel01}, or an evolved hydrogen main sequence star channel \\citep{pod03}. As shown by \\cite{del07} their subsequent evolution is affected by the thermal state of the donor star.  The observational investigation of these double white dwarf systems provides important insights into their nature and formation. For example, optical observations of the line spectrum from the system carry the combined physical and chemical properties of the accretor and the accretion disk while infrared observations of the spectrum may provide important constraints on the thermal state, chemical composition, and surface gravity of the donor. As a result, atmospheric studies of the donor in combination with knowledge of its distance can be used as an additional probe of the formation channel and evolutionary state of the system  \\citep{roel07}.  J0926+3624 offers a unique opportunity for probing the thermal properties of the donor star since its surface gravity, essential for calculating the synthetic spectrum, is known.  In addition, the orbital separation between the accretor and the donor, which determines the amount of irradiation onto the donor, is known from the component masses and orbital period.  In this Letter, we present the model composite spectra consisting of the non-LTE spectra of an irradiated donor, and a primary star as well as a multi-color blackbody spectrum of the accretion disk. The model parameters for the system are discussed in \\S 2. The formalism for estimating the contribution of the accretion disk is also discussed in this section. In \\S 3, a brief description of the numerical methods employed for calculating the non-LTE theoretical spectra of both the accretor and the donor stars is provided.  The results are described and discussed in \\S 4 followed by our conclusions in the last section. ",
        "conclusions": "Theoretical spectra were calculated for the two white dwarf components in the AM CVn system J0926+3624 using the non-LTE atmosphere code TLUSTY and SYNSPEC. A range of effective temperatures is adopted based on estimates derived from theoretical binary evolutionary considerations for a system formed in the white dwarf channel. Irradiation of the cool donor star by the hot primary star is included by assuming complete redistribution of energy over the entire atmosphere of the donor star. Observations indicate the presence of a spot whose contribution to the total flux is unknown. In the absence of a self-consistent non-steady state disk model, a multi-color blackbody spectrum of the disk is considered and a maximum disk size is adopted as inferred from observation. Comparison of the total flux derived by the present models places upper limits on the distance of the object as well as the contribution of the accretion disk to the total flux. The present findings emphasize that future infrared observations of the system would provide a good estimation of the contribution by the accretion disk to the total flux and hence important insight on the thermal properties of the accretion disk. We also predict emission lines in the infrared wavelength region, originating from the spectrum of the irradiated donor. The width of these lines are dependent on the effective temperature of both the donor and the accretor and hence high resolution spectrum may be needed to detect them. If, however, the donor is sufficiently cool, as suggested by the evolutionary models, the emission lines could be detected easily providing an important diagnostic probe for the donor."
    },
    "1107/1107.1819_arXiv.txt": {
        "abstract": "We observed the young pulsar J1357--6429 with the {\\it Chandra} and {\\it XMM-Newton} observatories. The pulsar spectrum fits well a combination of absorbed power-law model ($\\Gamma=1.7\\pm0.6$) and blackbody model ($kT=140^{+60}_{-40}$ eV, $R\\sim2$ km at the distance of 2.5 kpc). Strong pulsations with pulsed fraction of $42\\%\\pm5\\%$, apparently associated with the thermal component, were detected in 0.3--1.1 keV. Surprisingly, pulsed fraction at higher energies, 1.1--10 keV, appears to be smaller, $23\\%\\pm4\\%$.  The small emitting area  of the  thermal component either corresponds to a hotter fraction of the neutron star (NS) surface or indicates inapplicability of the simplistic blackbody description. The X-ray images also reveal a pulsar-wind nebula (PWN) with complex, asymmetric morphology comprised of a brighter, compact PWN surrounded by the fainter, much more extended PWN whose spectral slopes are $\\Gamma=1.3\\pm0.3$ and $\\Gamma=1.7\\pm0.2$, respectively. The extended PWN with the observed flux of $\\sim7.5\\times10^{-13}$ erg s$^{-1}$ cm$^{-2}$ is a factor of 10 more luminous then the compact PWN. The pulsar and its PWN are located close to the center of the extended TeV source HESS J1356--645, which strongly suggests that the VHE emission is powered by electrons injected by the pulsar long ago. The X-ray to TeV flux ratio, $\\sim0.1$, is similar to those of other relic PWNe. We found no other viable candidates to power the TeV source.  A region of diffuse radio emission, offset from the pulsar toward the center of the TeV source, could be synchrotron emission from the same relic PWN rather than from the supernova remnant. ",
        "introduction": "Thanks to their rich observational manifestations, young energetic pulsars are among the most attractive targets for the {\\sl Chandra} and {\\sl XMM-Newton} observatories. These manifestations include thermal emission from the neutron star (NS) surface, non-thermal emission from the pulsar magnetosphere, and a pulsar-wind nebula (PWN) produced by the interaction of the pulsar wind with the ambient medium. The magnetospheric and thermal emission can differ in their relative strengths and exhibit distinct spectra and pulse profiles. The magnetospheric emission usually has power-law (PL) spectrum and is strongly pulsed. Detailed studies of several bright young pulsars indicate a nonuniform NS surface temperature distribution with cool ($T_C \\lesssim 100$ eV) and hot ($T_H \\sim 0.1-0.3$ keV)  thermal components, which are believed to be emitted from the bulk of the NS surface and polar caps, respectively \\citep[e.g.][]{Pavlov2001,Zavlin2004,deLuca2005,Karga2005}. Since the sample of young pulsars with high quality X-ray spectra is still very small, further observations of young pulsars must be carried out to learn about the early epoch of NS cooling and test magnetospheric emission models.  In addition to X-ray emission from the NS surface and magnetosphere, a significant amounts of non-thermal emission is radiated from the PWN. PWNe display various shapes and structures, including (but not limited to) jets and tori, or cometary shaped tails \\citep{Gaensler2006}. Observations with high angular resolution are particularly useful since they resolve the extended PWN emission, seen downstream of the termination shock, from the magnetospheric non-thermal emission.  Comparing properties of these two components, one can learn about the physical processes responsible for an efficient but yet unknown particle acceleration mechanism energizing PWNe.  Although more than $\\sim$70 PWNe have been observed with {\\it Chandra} and {\\it XMM-Newton} to date \\citep[][hereafter KP08 and KP10]{Karga2008,Karga2010}, the existing models \\citep{Kennel1984,Komissarov2004,Swaluw2005,Bucciantini2005} still cannot explain some of the observed complex morphologies. Pulsars of similar ages and spin-down luminosities can have PWNe with quite different properties. As of now,  it is not clear whether these differences can be solely attributed to different properties of the surrounding interstellar medium (ISM) or different intrinsic properties of pulsars also play an important role. A larger sample of well-resolved X-ray PWNe associated with pulsars having different intrinsic properties and residing in different environments  is needed to further advance our understanding of pulsar winds.  The young pulsar PSR J1357--6429 (hereafter J1357; $P$ = 166 ms, $\\dot{E}$ = 3.1 $\\times$ 10$^{36}$ erg s$^{-1}$, surface magnetic field $B$ = 7.8 $\\times$ 10$^{12}$ G) was discovered in the Parkes Multibeam Pulsar Survey by \\citet{Camilo2004}. The spin-down age, $\\tau \\equiv P/(2 \\dot{P})$ = 7.3 kyr, indicates that J1357 is one of the youngest pulsars known. The distance of 2.5 kpc, estimated from  the pulsar's dispersion measure (DM = 127.2 cm$^{-3}$ pc) and the Galactic free electron density model \\citep{Cordes2002}, implies a spin--down flux $\\dot{E}/(4 \\pi d^2)$ = 1.2 $\\times$ 10$^{-9}$ erg s$^{-1}$ cm$^{-2}$.  The pulsar is located within the TeV source HESS J1356--645 \\citep[hereafter HESS J1356;][]{Renaud2008,Abramowski2011}. The extended TeV emission is detected up to $12'$ from the peak of the TeV surface brightness, located at  R.A.$ =13^{\\rm h}56^{\\rm m}$, decl.$=-64^{\\circ}30'$. Given the relatively small angular separation, $\\sim 7'$, between J1357 and the center of the HESS source, and the TeV luminosity, $L_{\\rm TeV}\\sim 6 \\times 10^{33}$ ergs s$^{-1} \\sim 0.002\\dot{E}$ (at the 2.5 kpc distance), it seems plausible that the HESS source could be powered by this pulsar. The TeV spectrum of HESS J1356 fits PL with the photon index $\\Gamma_{\\rm TeV}\\approx2.2$, without a cutoff at high energies. A possible detection of J1357 with {\\sl AGILE}  at lower $\\gamma$-ray energies (100 MeV -- 10 GeV) was reported by \\citet{Pellizzoni2009}\\footnote{After our paper was submitted, \\citet{Lemoine2011} reported results of observations of J1357 with {\\sl Fermi} LAT. They detected $\\gamma$-ray pulsations and measured the photon index $\\Gamma_{\\rm GeV}\\approx 1.5$, found an exponential cut-off at $\\approx 0.8$ GeV, and estimated the luminosity $L_{\\rm GeV} \\approx 2\\times 10^{34}$ erg s$^{-1}$, for $E>0.1$ GeV.}. No optical counterpart to the pulsar or its PWN has been detected yet \\citep{Mignani2011}.  At radio frequencies, \\citet{Duncan1997} reported the extended emission from  the SNR candidate G309.8--2.6 , which is spatially coincident with HESS J1356. The extent of the radio emission appears to be smaller than that of the TeV source. J1357 has been previously observed with {\\sl XMM-Newton} (2005 August 5; 15 ks; PI F.\\ Camilo) and {\\sl Chandra} HRC-S (2005 November 18 and 19; 16 and 17 ks; PI M.\\ Mendez, observer L.\\ Kuiper). The analysis of those data was reported by \\citet{Zavlin2007} and \\citet{Esposito2007}\\footnote{It should be noted that first report on the {\\sl Chandra} results was published in ``Chandra News'' (March 2007, Issue number 14, pages 12-13 (see \\url{http://cxc.harvard.edu/newsletters/news\\_14/newsletter14.html}) by R.\\ Kraft and A.\\ Kenter. In particular, evidence for the presence of extended emission (PWN) and the detection of pulsed signal (pulsed fraction $40\\%\\pm 12\\%$) were reported in that publication.}. These authors found that the pulsar spectrum likely consists of thermal and nonthermal components. \\citet{Zavlin2007} found pulsations with the radio pulsar period and pulsed fraction of $63\\%\\pm 15\\%$, while \\citet{Esposito2007} reported only an upper limit of 30\\% on modulation amplitude. \\citet{Zavlin2007} also reported a tail-like extended emission in the HRC-S image, with a luminosity of $2.5\\times 10^{31}$ erg s$^{-1}$ (for $d=2.5$ kpc), while \\citet{Esposito2007} concluded that the surface brightness distribution around the pulsar ``is consistent with that from a point source'' and estimated a 2--10 keV upper limit of $3\\times 10^{31}$ erg s$^{-1}$ on the PWN luminosity. To resolve these controversies, reliably disentangle various emission components, and study the relation between the pulsar and HESS~J1356, we have carried out deeper observations with the {\\sl Chandra} ACIS and {\\sl XMM-Newton} EPIC detectors.  The details of the observations and data analysis are presented in Section 2. We discuss  possible interpretations of the PWN morphology, describe inferences from the pulsar spectrum,  and speculate on the nature of HESS~J1356 and its relation to the other sources in the field in Section 3. ",
        "conclusions": "Our {\\sl Chandra} ACIS and {\\sl XMM-Newton } EPIC observations have shown that the PWN of J1357 is much more extended and luminous  than it could be seen from the earlier observations. We have also found that the thermal component of pulsar emission is strongly pulsed, while the pulsation of the nonthermal component, which dominates at higher energies, are much weaker. Below we discuss the physical implications of our findings, compare the J1357 PSR/PWN properties to those of other similar pulsars observed in X-rays, and  investigate the relation between J1357, HESS~J1356, and putative host SNR G309.8--2.6.  \\subsection{Pulsar}  The pulsar spectrum cannot be satisfactory fitted by simple one-component models, such as absorbed PL or BB (see \\S2.2.1).  The two-component BB+PL model fits the spectrum much better, although some residuals are still present in the combined fits. The best-fit photon index for the PL component, $\\Gamma\\simeq1.7\\pm0.6$, is marginally larger than $\\Gamma\\simeq1.3\\pm0.3$ of the compact PWN (region PWN1). The non-thermal luminosity is $L_{\\rm nonth}\\approx5\\times10^{31}$ erg s$^{-1}$ in the energy range 0.3--8 keV. The PL component may partly come from an unresolved PWN. This could explain the weak pulsations above 1 keV (see \\S2.2.3), in contrast to, for instance, the Vela pulsar (which is well resolved from the surrounding PWN), whose pulsed fraction grows with energy reaching 62\\% above 1.8 keV \\citep{Sanwal2002}. Our timing analysis limits the luminosity of the unresolved PWN to $\\lesssim 77\\%$ of the point source luminosity in 1.1-10 keV.   Unlike the hard non-thermal component, the soft thermal component should not be strongly contaminated by the PWN emission since it clearly dominates at $\\lesssim1$ keV. Therefore, the best-fit BB temperature, $T\\approx1.9-2.7$ MK, radius, $R\\sim 0.3-1.1$ km, and bolometric luminosity, $L_{\\rm bol}\\sim 2\\times10^{31}$ erg s$^{-1}$, are reasonably accurately determined (see Table~\\ref{tab2}). The values are similar to those of other Vela-like pulsars. Although the temperature is a factor of two higher than expected from the standard NS cooling, and the emitting area radius is a factor of ten smaller then the classical NS radius, the discrepancies can be attributed to the use of the simplistic BB model. Indeed, the fit with the hydrogen atmosphere models \\citep[e.g.,][]{Pavlov1995}, described in Section 2.2.1, gives a lower temperature, $T_{\\rm eff}^\\infty = 0.70$--0.78 MK at $R^\\infty = 13$ km, at the same bolometric luminosity. The measured bolometric luminosity is a factor of $\\sim 3$ lower than predicted by the theoretical cooling curve for low-mass NSs ($M_{\\rm NS} = 1.3 M_\\odot$) without superfluidity; it is consistent with the cooling curves for heavier NSs (e.g., $M_{\\rm NS}=1.5$--$1.6\\, M_\\odot$) for some superfluidity models \\citep{Yakovlev2004}.  Very surprising is the high ($\\approx40$\\%) pulsed fraction associated with the thermal component. In general, one should not expect such strong pulsations from a non-uniformly heated NS surface with dipolar magnetic field for a blackbody-like angular distribution. However, there have been recent reports of even stronger thermal pulsations in strongly magnetized young pulsars, such as PSR J1119--6127 ($74\\%\\pm14\\%$; Gonzalez et al.\\  2005) and PSR J1718--3718 ($52\\%\\pm13\\%$; Zhu et al. 2011). This suggests that the strong pulsations are caused by the combined effect of the anisotropic temperature distribution and strong anisotropy of emitted radiation in high magnetic fields \\citep{Pavlov1994}. An alternative possibility is that part of emission interpreted as thermal from our BB+PL and NSA+PL fits is in fact an additional, softer non-thermal component, which is masked at higher energies by the harder, unresolved PWN emission.  This unaccounted component might be the reason for the remaining residuals of the PL+BB and NSA+PL fits. However, higher quality data are needed to test this hypothesis.   \\subsection{PWN}  At the plausible distance of 2.5 kpc, the unabsorbed X-ray luminosity of the compact PWN, $L_{\\rm pwn1}\\approx 5 \\times 10^{31}$ ergs s$^{-1}$ in the 0.3-8 keV band, corresponds to the X-ray efficiency, $\\eta_{\\rm pwn1}\\equiv L_{\\rm pwn1}/\\dot{E} \\sim 1.7\\times10^{-5}$, lower than those of most Vela-like pulsars (see Kargaltsev et al.\\ 2007; KP08). This suggests  that either the distance is underestimated or, in addition to $\\dot{E}$ and $\\tau$, $\\eta_{\\rm pwn1}$ strongly depends on other factors (e.g., the pulsar's speed and the angle between the spin and magnetic axes). Note, however, that the efficiency of the entire detected PWN is higher by a factor of 15, $\\eta_{\\rm pwn}\\sim 2.6\\times10^{-4}$.  The spectral slope of the compact PWN, $\\Gamma_{\\rm pwn1} =1.3\\pm0.3$, is similar to those of PWNe around Vela-like pulsars \\citep[listed in Table 2 of][]{Karga2007c}. The PWN spectrum apparently softens with the distance from the pulsar, becoming  $\\Gamma\\approx1.7-1.8$ for the narrow feature (PWN2) and the large-scale PWN, which suggests synchrotron cooling of the outflow. The degree of spectral softening is similar to that seen for the extended tail of  PSR J1509--5850, where the photon index changes from $\\Gamma=1.8\\pm0.3$ for PWN in the vicinity of the pulsar to $2.4\\pm0.4$ for the extended tail \\citep{Karga2008a}.  Our X-ray images reveal the morphology of the PWN on different  angular scales. The observed  morphology does not easily  fit in either bow-shock or torus-jet category  (see KP08).  The bright compact PWN1 extends up to $\\simeq 10''$ from the pulsar and  has rather amorphous, somewhat asymmetric  morphology.  The diffuse emission is noticeably brighter northeast of the pulsar compared to the opposite side, although some emission is still clearly discernible within the PWN1 region southwest of the pulsar. We find no evidence of  structures that could be interpreted as a torus associated with the termination shock (TS) in the pulsar wind, which may suggest that the TS occurs at an angular distance of several arcseconds or less, corresponding to $r_s\\lesssim4\\times10^{16}d_{2.5}$ cm (versus $1\\times10^{17}$ cm in the Vela PWN; Helfand et al.\\ 2001). If the pulsar is moving subsonically inside the hot SNR medium, we can use the above estimate to obtain a lower limit on the ambient pressure at $p_{\\rm amb}\\sim \\dot{E}  (4\\pi c r_s^2)^{-1}\\gtrsim5\\times 10^{-9}(r_s/4\\times10^{16}~{\\rm cm})^{-2}$ dyn cm$^{-2}$, which is close to the maximum pressure one can expect in a 10 kyr old SNR \\citep{Karga2009, Bamba2010}. The high pressure required, and the asymmetry of the compact PWN, elongated approximately along the northeast-southwest direction, might be indicative of the TS being also asymmetric, i.e., being farther away from the pulsar on the northeast side than on the southwest side. Such situation may occur if the pulsar is moving fast through the surrounding medium in the southwest direction, and the ram pressure of the medium compresses the PWN in front of the pulsar. If the pulsar's speed exceeds the sound speed in the medium, $c_s=(5 kT/3 \\mu m_{\\rm H})^{1/2}=12\\, \\mu^{-1/2} T_4^{1/2}$ km s$^{-1}$, a bowshock is formed. The appearance of a bowshock PWN depends on the Mach number, uniformity of the medium, and intrinsic anisotropy of the pulsar wind.  X-ray PWNe around supersonically moving pulsars often exhibit a bright ``bullet'' in the vicinity of the pulsar and a much more extended faint tail in the direction opposite to that of the pulsar motion; however, more complex structures have been seen in some cases (e.g., the Guitar and Eel PWNe associated with  B2224+65 and J1826--1256, respectively; \\citet{Johnson2010, Roberts2007}).  The adjacent PWN2 region shown in Figure 1 encompasses the fainter narrow feature attached to the compact PWN. Interestingly, this narrow feature bends sharply at about $23''$ from the pulsar. Although there is no doubt that the feature is part of the PWN, the origin of the feature is not clear. For instance, it could be a pulsar jet. Indeed, some of pulsar jets are known to show extreme bending (e.g., the outer jet of the Vela pulsar; \\citet{Pavlov2003}). The apparent lack of a similarly looking counter-jet should not be surprising since the two jets can differ significantly both in shape and surface  brightness due to the Doppler boost and/or proper motion effects  (e.g., \\citet{Pavlov2003}). On the other hand, the ratio of the compact PWN luminosity to that of the putative jet, $L_{\\rm pwn1}/L_{\\rm pwn2}\\sim5$, is noticeably larger than that in the Vela and Crab PWNe but comparable to that of the PWN around PSR B1706--44 \\citep{Romani2005}. If this interpretation is taken at the face value, one could expect the pulsar to be moving northeast (in general, jets tend to be co-aligned with the pulsar's direction of motion--cf.\\ the Vela pulsar; \\citet{Pavlov2003}).   Alternatively, the narrow feature could be a part (e.g., an inner channel) of the pulsar tail (see, e.g., \\citet{Karga2008} and references therein). In this case the pulsar would be moving in the opposite (southwest) direction. This interpretation is  in line with the above-described asymmetry of the compact PWN whose brighter part could, in this case, indeed be the ``bullet'' associated with the termination shock deformed by the ram pressure of the oncoming ISM.  However, the sharp bending and the lack of a fainter, much more extended tail (cf.\\ J1509--5850 tail; \\citet{Karga2008}), and the emission west-southwest of the compact PWN, are at odds with this interpretation\\footnote{The diffuse emission seen in the PMN 4.85 GHz image north of the pulsar (panel e in Fig.~6) could, in principle, be the radio counterpart of the extended pulsar tail. Deeper radio observations are required to test this hypothesis.}.  Finally, the narrow feature could be a part of the bow associated with the forward shock (see e.g., the Geminga PWN; Pavlov et al.\\ 2006, 2010) with the pulsar moving southeast.  In this case the observed large-scale emission can also be associated with the asymmetric bowshock similar to that seen around  PSR J1826--1256 \\citep{Roberts2007}. Somewhat puzzling may be the lack of any X-ray emission within the bow interior; however, this is also the case in the Guitar and Eel bowshocks. To conclude, with the data in hand we cannot univocally interpret the morphology of the X-ray PWN. Measuring the pulsar's proper motion should provide decisive  information allowing one to discriminate between the above interpretations. Assuming a typical pulsar speed of 400 km s$^{-1}$ \\citep{Hobbs2005}, this should be possible to accomplish in several years with VLBI.  \\begin{table*} \\footnotesize \\center \\caption{Properties of candidate TeV plerions and their parent pulsars.\\label{tab4}} \\vspace{0.5cm} \\setlength{\\tabcolsep}{1mm} \\begin{tabular}{lccccccccccc} \\hline\\hline HESS ID &  $f_{\\gamma}$\\tablenotemark{a} & TeV size &  Offset\\tablenotemark{b}  & PSR & $\\dot{E}_{36}$\\tablenotemark{c} & $\\tau$ & Dist. & $L_{\\gamma}/\\dot{E}$ &  $L_{X}/L_{\\gamma}$ &  $\\mathcal{B}_{\\rm acis}$\\tablenotemark{d}   \\\\ &   & pc &  &  &  & kyr & kpc & & &\\\\\\hline J1809$-$193 & 0.3 & $20$ & $8'$ &  J1809--1917 & 1.8  & 50 & 3.5 & 1\\% & $2\\%$ & 8  \\\\ Vela X  & 0.75 & $5$ & $30'$ & J0835--4510 & 6.9 & 11 & 0.3 & 0.002\\% & $90\\%$ & -- \\\\ J1825$-$137 & 0.48 & $70$ & $10'$ & J1826--1334 & 2.8 & 21 & 3.5 & $4\\%$ & $0.3\\%$ & 7.5  \\\\ J1356$-$645 & 0.79 & $40$ & $8'$  & J1357--6429 & 3.1 & 7.3 & 2.5 & $0.2\\%$ & $10\\%$ & 19 \\\\ \\hline\\hline \\end{tabular} \\tablenotetext{\\rm a}{~Unabsorbed $\\gamma$-ray flux (1--10 TeV) in units of $10^{-11}$ ergs s$^{-1}$ cm$^{-2}$.} \\tablenotetext{\\rm b}{~Offset of the TeV source center from the pulsar.} \\tablenotetext{\\rm c}{~Pulsar spin-down power, $\\dot{E}$, in units of  $10^{36}$ ergs s$^{-1}$.} \\tablenotetext{\\rm d}{~Average surface brightness  measured from {\\sl Chandra} ACIS images, in counts ks$^{-1}$ arcmin$^{-2}$.} \\end{table*}  \\subsection{The Nature of the VHE source}  Given the lack of other promising counterparts and the location of J1357 within the HESS~J1356 extent, it seems very plausible that the two should be related in some way.  Yet the nature of HESS~J1356 has not been firmly  established.  The TeV emission could be attributed to either the host SNR or to the relic PWN of J1357.  To date, several well-known SNRs have been detected in the TeV band \\citep[e.g.,][]{Aha2004}. In the cases of resolved SNRs, the TeV emission was associated with the SNR shell, and the sizes of the SNRs in TeV images were similar to those of the non-thermal radio shells indicating that both emission components are powered by particles accelerated in the forward shock.  Although no shell is seen in the radio images of HESS~J1356, the radio emission is mainly seen  near the center of HESS~J1356 (see Fig.~\\ref{fig6}). Thus, in this case the TeV emission does not appear to be associated with the nonthermal shell, and the shell itself is not seen in the radio.  An alternative, more plausible, interpretation of the TeV and radio emission could be a relic PWN. Many extended TeV sources  neighbor young Vela-like pulsars, sometimes offset up to $10'$--$20'$ from the center of TeV emission (see KP10 for recent review). Furthermore, recent X-ray observations of the Vela pulsar region  (Vela X; Mori et al.\\ 2008),  PSR~J1826--1334 (Gaensler et al.\\ 2003; Pavlov et al.\\ 2008) and PSR~J1809--1917 \\citep{Karga2007} have provided convincing evidence that the TeV sources are connected to the parent pulsars through faint asymmetric X-ray nebulae\\footnote{There remains a large fraction of relic PWN candidates where only bright compact PWN is seen X-rays while the extended component is not seen perhaps due to is faintness and/or rapid cooling of the outflow.} (see Table~\\ref{tab4}). Indeed, the {\\sl Chandra} observations have demonstrated that in the above three examples the PWNe consist of compact (0.1--0.8 pc) bright cores and  fainter asymmetric, more extended (2--8 pc) components. The  offsets and the asymmetries of the X-ray PWNe could be created by the reverse SNR shock that had propagated through the nonhomogeneous SNR interior and reached one side of the PWN sooner than the other side \\citep{Blondin2001}. This scenario could also account for the similarly asymmetric, offset TeV emission (e.g., de Jager \\& Djannati-Ata\\\"{\\i} 2009).  However, the physical origin of the TeV emission still remains under debate. It can be produced via the Inverse Compton Scattering (ICS) of the relativistic pulsar-wind electrons on photons from the omnipresent cosmic microwave background (CMB), galactic IR background, and IR photons from local star-forming regions and warm dust clouds. Alternatively, the TeV photons can be produced as a result of $\\pi^{0}\\rightarrow\\gamma + \\gamma$ decay, with $\\pi^{0}$ being produced when the relativistic protons from the pulsar wind interact with the ambient matter (e.g., Horns et al.\\ 2007). In the case of HESS~J1356, the offset is well within the range observed in other relic PWN candidates (see KP10). The ratio of the TeV luminosity to the pulsar's spin-down power, $L_{\\rm TeV}/\\dot{E} \\sim 0.002$, and the ratio of X-ray to $\\gamma$-ray (1--10 TeV) PWN luminosities, $L_X/L_{\\rm TeV}\\sim 0.1$, are similar to those of other relic plerions.  Alternatively, the relic PWN could be left behind the fast moving pulsar if the latter is moving in the northeast direction (see above). At a typical speed of 400 km s$^{-1}$, the pulsar would move by $5\\farcm 6$ in 10 kyrs, which is consistent with the offset between the pulsar and the center of HESS~J1356. So far, there have been no reports of VHE emission from pulsar tails. It would be interesting to establish the presence of the TeV emission from an extended  pulsar tail because this would help to break the degeneracy in interpreting the nature of the VHE emission from the crushed PWNe, for which both leptonic (Inverse Compton) and hadronic ($\\pi^0$ decay) TeV emission mechanisms are currently being debated. If the $\\pi^0$ decay is the dominant process in the crushed PWNe, we would expect pulsar tails to be significantly fainter in the TeV  because high-speed pulsars move in low density media outside their host SNRs. On the other hand,  in the leptonic scenario, crushed plerions and pulsar tails should have comparable TeV luminosities. Regardless of the nature of the relic PWN and the mechanism responsible for the VHE emission, one can expect that the observed extended radio emission within HESS~J1356 may also be produced by the pulsar wind if the electron SED extends to sufficiently low energies.  Better quality radio images, more complete multiwavelength spectrum, and robust multizone pulsar wind models are required to assess the origin of the radio emission."
    },
    "1107/1107.3507_arXiv.txt": {
        "abstract": "New observations show that the lightcurve of Kuiper belt contact binary (139775) 2001$\\,$QG$_{298}$ has changed substantially since the first observations in 2003. The 2010 lightcurve has a peak-to-peak photometric of range $\\Delta m_{2010}=0.7\\pm0.1$ mag, significantly lower than in 2003, $\\Delta m_{2003}=1.14\\pm0.04$ mag. This change is most simply interpreted if 2001$\\,$QG$_{298}$ has an obliquity near $90\\degr$. The observed decrease in $\\Delta m$ is caused by a change in viewing geometry, from equator-on in 2003 to  nearly 16$\\degr$ (the orbital angular distance covered by the object between the observations) off the equator in 2010. The 2003 and 2010 lightcurves have the same rotation period and appear in phase when shifted by an integer number of full rotations, also consistent with high obliquity. Based on the new 2010 lightcurve data, we find that 2001$\\,$QG$_{298}$ has an obliquity $\\varepsilon=90\\degr\\pm30\\degr$.  Current estimates of the intrinsic fraction of contact binaries in the Kuiper belt are debiased assuming that these objects have randomly oriented spins. If, as 2001$\\,$QG$_{298}$, KBO contact binaries tend to have large obliquities, a larger correction is required. As a result, the abundance of contact binaries may be larger than previously believed. ",
        "introduction": "Located approximately between 30 and 50 AU from the Sun, the Kuiper belt represents the outer frontier of the currently observable solar system.  The belt is estimated to hold roughly forty thousand objects larger than 100 km in diameter \\citep{1995AJ....109.1867Jewitt,2001AJ....122..457T,2004AJ....128.1364B,2008AJ....136...83Fuentes,2008Icar..195..827Fraser}. More than one thousand objects have been detected but only six hundred have multi-opposition observations and well-determined orbits. The Kuiper belt objects (KBOs) are remnants of outer solar system planetesimals and their study may shed light on the complex process that led to the formation of planets. Our understanding of the physical properties of KBOs remains rudimentary as most are too faint for detailed investigation.  Many KBOs are binary.  Among the dynamically cold, ``classical'' population (low eccentricity and inclination KBOs located between the 3:2 and the 2:1 mean-motion resonances with Neptune, at 39.4 and 47.8 AU) the binary frequency is inferred to be $\\sim$20\\% whereas in other dynamical subsets the fraction is lower, at 5 to 10\\% \\citep{2006AJ....131.1142Stephens,2008Icar..194..758Noll}. Binaries probably formed early on \\citep{2002Icar..160..212W}, before the Kuiper belt lost more than 99\\% of its original mass \\citep{1998AJ....115.2136K,2009MNRAS.399..385Booth}, because the current-day number density of KBOs is too low for frequent pair encounters. It is also possible that planetesimals formed in clusters of two or more bodies and that the binaries we see today are direct remnants of that process \\citep{2010AJ....140..785Nesvorny}. Thus, binaries are particularly useful probes of the conditions under which planetesimals formed. For instance, the symmetry in surface properties of KBO pairs argues for a chemically homogeneous formation environment \\citep{2009Icar..200..292Benecchi}. But the current binary abundance also bears record of the effects of collisional and dynamical processes that have eroded the original population \\citep{2004Icar..168..409P}. It is highly desirable to find ways to isolate the effects of formation, evolution, and erosion.  The relative frequency of binaries within different regions of the Kuiper belt, and their distributions of orbital properties and relative mass of the components are quantities that should be sought to help clarify the formation and evolution of binaries \\citep{2008ApJ...686..741Schlichting,2011ApJ...730..132Murray-Clay,2011AJ....141..159Nesvorny}.  Most known binaries are resolved \\citep[physical separations $>$1500 km;][]{2002AJ....124.3424N,2006AJ....131.1142Stephens,2006ApJ...643..556Stansberry,2011arXiv1103.2751Grundy}. The exception is the $500\\,\\mathrm{km}\\times175\\,\\mathrm{km}$ contact binary (139775) 2001$\\,$QG$_{298}$, identified from analysis of its rotational lightcurve \\cite[Figure \\ref{Fig.One};][hereafter SJ04]{2004AJ....127.3023S}. Using time-resolved measurements taken mainly during 2003, the authors found that 2001$\\,$QG$_{298}$ displayed extreme photometric variability ($\\Delta m\\sim1.2$ mag) and a relatively slow rotation period ($P\\sim13.77$ hr). The large photometric range, caused by the eclipsing nature of the binary, exceeds the $\\Delta m=0.75$ mag produced by two spheres in contact (axis ratio $b/a=0.5$) because the contact binary components are tidally elongated along the line connecting the centers \\citep{1980Icar...44..807W,1984A&A...140..265L}. The large $\\Delta m$ found for 2001$\\,$QG$_{298}$ also implies that the system was observed almost equatorially in 2003. Models of 2001$\\,$QG$_{298}$'s lightcurve based on figures of hydrostatic equilibrium lend strong support to these assertions \\citep{2004PASJ...56.1099T,2007AJ....133.1393L,2010ApJ...719.1602Gnat} and allow the bulk density of 2001$\\,$QG$_{298}$ to be estimated at $\\rho=0.59_{-0.05}^{+0.14}$ g$\\,$cm$^{-3}$. This surprisingly low density implies that 2001$\\,$QG$_{298}$ is mostly icy in composition and is significantly porous.  The example of 2001$\\,$QG$_{298}$ suggests that between 10\\% and 20\\% of KBOs could be contact binaries, depending on surface scattering properties \\citep[SJ04;][]{2007AJ....133.1393L}. This high fraction is consistent with the observation that the frequency of resolved binaries increases dramatically with decreasing orbital separation \\citep{2006ApJ...643L..57Kern}.  However, despite their high abundance, contact binary KBOs remain elusive. The reason is that these objects are unresolved and can only be identified if they happen to be nearly equator-on. The consequence is that as many as 85\\% of contact binaries may go unnoticed due to unfavorable observing geometry. High phase angle observations that have been shown to increase the chance of detecting contact binaries \\citep{2008ApJ...672L..57Lacerda} can not be applied to the distant KBOs as their phase angles remain below 2\\degr\\ when observed from Earth. Interestingly, the Jupiter Trojans also contain a fairly large fraction of contact binaries, comparable to that found for KBOs. In addition to the well-known case of (624) Hektor \\citep{1978Icar...36..353H,1994A&A...281..269Detal,2001Icar..153..348Cruikshank,2007AJ....133.1393L}, two other contact binaries have been identified in a survey of 114 Trojans \\citep{2007AJ....134.1133Mann}. When corrected by the probability of observing these objects close to equator-on, \\citet{2007AJ....134.1133Mann} find an intrinsic contact-binary fraction of at least 6\\% to 10\\%. The apparently high abundance of contact binaries and the prospect of measuring their densities makes these primitive objects particularly interesting targets of study.  QG$_{298}$ has travelled a significant angular distance ($\\sim16\\degr$) in its heliocentric orbit since the first observations in 2003. In this paper we compare new 2010 lightcurve data with the SJ04 observations and make inferences on the obliquity of 2001$\\,$QG$_{298}$. In Section 2 we describe our observations, in Section 3 we report our findings on the lightcurve and obliquity of 2001$\\,$QG$_{298}$, in Section 4 we discuss the implications of our findings for the abundance of contact binaries, and in Section 5 we present our conclusions. ",
        "conclusions": "We have measured the visible lightcurve of Kuiper belt object 2001$\\,$QG$_{298}$ and compared it to similar data collected in 2003. Our findings can be summarised as follows:  \\begin{enumerate}  \\item The 2010 lightcurve of 2001$\\,$QG$_{298}$ has a peak-to-peak range of $\\Delta m_{2010}=0.7\\pm0.1$ mag, significantly lower than the photometric range in 2003, $\\Delta m_{2003}=1.14\\pm0.04$ mag. The 2003 and 2010 lightcurves appear in phase when shifted by an integer number of full rotations.  \\item The change between the 2003 and 2010 lightcurves is most simply explained if 2001$\\,$QG$_{298}$ possesses a large obliquity, $\\varepsilon=90\\degr\\pm30\\degr$. In that case, the lightcurve photometric range should continue to decrease, reaching a minimum of $\\Delta m\\sim0.0-0.1$ mag in 2049.  \\item Current estimates of the fraction of contact binaries in the Kuiper belt assume that these objects have randomly oriented spins. If, as 2001$\\,$QG$_{298}$, contact binaries tend to have large obliquities their abundance may be larger than previously believed.  \\end{enumerate}"
    },
    "1107/1107.3131.txt": {
        "abstract": "Fractionally charged massive particles (FCHAMPs) appear in extensions of the standard model, especially those with superstring constructions.  The lightest FCHAMP would be absolutely stable and any of them produced during the early evolution of the Universe would be present today.  The thermal production, annihilation and, survival of an FCHAMP, a lepton \\L~with electroweak (i.e., $U(1)_Y$) but no strong interactions, of mass \\ml~and charge \\ql~(in units of the charge on the electron) are explored.  The FCHAMP relic abundance is determined by the total annihilation cross section which depends on \\ml, \\ql~and on the available annihilation channels.  Since massive ($m_{L} \\ga 1$~GeV) charged particles ($Q_{L} \\ga 0.01$) behave like baryons (heavy ions), primordial nucleosynthesis and the cosmic background radiation temperature anisotropies limit the FCHAMP relic density.  Requiring that $\\Omega_{L} \\la \\Omega_{\\rm B}/5$ leads to a constraint on the \\ql~-- \\ml~relation.  Further constraints on \\ql~and \\ml~are provided by the invisible width of the $Z$ ($Q_L > 0.16$ for $m_{L} \\leq M_{Z}/2$) and by accelerator searches for massive, charged particles.  Our key result is to exploit the fact that in the early Universe, after \\Lpm~freeze-out but prior to \\epm recombination, the negatively charged $L^{-}$ will combine with the more abundant alpha particles and protons to form tightly bound, positively charged states (fractionally charged heavy ions).  The Coulomb barriers between these positively charged $L^{-}\\alpha$ and $L^{-}p$ ($L^{-}pp$, $L^{-}\\alpha\\alpha$, ...) bound states and the free $L^{+}$ suppresses late time FCHAMP annihilation in the interstellar medium (ISM) of the Galaxy and on Earth, limiting significantly the late-time reduction of the FCHAMP abundance compared to its relic value.  The surviving FCHAMP abundance on Earth is orders of magnitude higher than the limits from terrestrial searches for fractionally charged particles, closing the window on FCHAMPs with $Q_{L} \\ga 0.01$.  However, as \\ql~approaches an integer (\\eg $|Q_{L} - n| \\la 0.25$) these searches become increasingly insensitive, leaving some potentially unconstrained ``islands\" in the \\ql~-- \\ml~plane which may be explored by searching for FCHAMPs in the cosmic rays. ",
        "introduction": "\\label{intro}  We consider the cosmological and astrophysical viability of a fractionally charged massive particle, an FCHAMP, the lightest of which is stable on cosmological timescales.  While there is an extensive literature on integer charged CHAMPs and neutraCHAMPs and their neutral counterparts in CHAMP-proton and CHAMP-alpha particle bound states (see, \\eg \\cite{khlopov,derujula,dimopoulos,gould,davidson}), in very few of these papers were fractionally charged FCHAMPs explored in any detail.  Many of these previous constraints on CHAMPs do not apply to FCHAMPs.  Our study is restricted to a fractionally charged lepton $L$ (an  \\Lpm~particle -- antiparticle pair), an $SU(3)$ and $SU(2)$ singlet with weak hypercharge $Y=Q_L$, which couples to the photon and the $Z$.  This choice fixes the total \\Lpm~annihilation cross section, determining the relic abundance of \\Lpm~pairs in the interstellar medium (ISM) of the Galaxy and on Earth.  Requiring that the relic abundance (mass fraction) be bounded limits a combination of the mass, $m_{L}$, and charge, $Q_{L}$, providing a lower bound to \\ql~as a function of \\ml~\\cite{wolfram,steigman1979,goldberg}.  During the post-annihilation freeze out evolution of the Universe, prior to recombination, the negatively charged $L^{-}$ will combine with alpha partcles and protons to form tightly bound, positively charged, heavy ions such as $L\\alpha$, $Lp$, $Lpp$, $L\\alpha\\alpha$.  The net positive charges of these exotic heavy ions suppresses late time annihilation of the bound $L^{-}$ with the free $L^{+}$.  This suppression leads to a terrestrial ratio of FCHAMPs to baryons which is so high that terrestrial searches for fractionally charged particles~\\cite{Perl:2009zz} challenge their very existence.  In \\S\\ref{motivation} the particle physics motivations for the existence of FCHAMPs are outlined along with general comments on existing experimental limits.  In \\S\\ref{relic} the relic FCHAMP abundance is calculated and limits are set to $m_{L}$ and $Q_{L}$.  Constraints on \\ql~and \\ml~from collider searches and from the invisible width of the $Z$ are analyzed in \\S\\ref{acc}, and those from terrestrial searches for fractionally charged particles in \\S\\ref{earth}.  The post-annihilation freeze out capture of negatively charged FCHAMPs by protons and alpha particles is described and the ratio of $Lp$ and $L\\alpha$ to baryons in the present Universe is estimated in \\S\\ref{capture}.  In \\S\\ref{discussion} we compare our results to the bounds from the terrestrial searches to decide if they exclude FCHAMPs -- or not.  Our conclusions are summarized in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  FCHAMPs with $Q_{L} \\ga 0.01$ and $m_{L} \\ga$~few GeV behave like heavy baryons~\\cite{haibo}, leading to an interesting constraint on their relic mass density~\\cite{rubtsov}, $\\Omega_{L} \\la \\Omega_{\\rm B}/5$ ($\\Omega_{L}h^{2} \\la 0.0044$).  However, if this constraint were saturated the relic FCHAMP abundance on Earth would be orders of magnitude larger than the limits set by the negative results of searches for terrestrial, fractionally charged particles.  But, if the relic terrestrial FCHAMPs consisted of ``free\" \\Lpm~pairs, renewed, Sommerfeld enhanced, annihilation would have occurred over the past 4.5 Gyr in the high density, low temperature environment of the Earth, reducing the relic abundance and resulting in a range of \\ql~-- \\ml~values ($Q_{L} \\ga 0.2$, $m_{L} \\ga M_{Z}/2$; see Figure~\\ref{qvsmall2}) consistent with the upper bounds inferred from the terrestrial searches.  However, as we have noted (see \\S\\ref{capture}), in the early Universe after \\Lpm~freeze out, negatively charged FCHAMPs will be captured into positively charged bound states with alpha particles and protons, suppressing any further, late time annihilation.  This would appear to close the window on {\\bf any} FCHAMPs with $0.01 \\la Q_{L} <2~(< 4)$.  Although a ``natural\" upper bound for the range of the electric charge of an FCHAMP might be $Q_{L} < 1$, we have extended the analysis here to $Q_{L} < 2$ (or, to $Q_{L} < 4$ by including $L\\alpha\\alpha$). We expect that larger values of $Q_L$ would also be excluded by the negative results of searches for fractionally-charged particles on Earth because of bound states with multiple $\\alpha$'s and/or heavier nuclei.  In that case, Eq.~\\ref{abundance} would suggest that only enormous values of $Q_L (> 10^4-10^5$) would be allowed.  We have not considered the very large $Q_L$ scenario here because it would likely require modifying BBN and, furthermore, perturbative calculations would break down for $\\alpha Q_L^2/4 \\pi =\\mathcal{O}(1)$.  Note that for \\ml~$\\la 300$~GeV, the D0 constraint~\\cite{Acosta:2002ju} suggests that $Q_{L,max} \\la 10.3$, resulting in a {\\bf lower} bound to the relic FCHAMP to baryon ratio, $n_{L}/n_{\\rm B} \\ga 7\\times 10^{-7}$.  Having argued that the extremely low upper bound from terrestrial searches for FCHAMPs has closed the window on such particles, we then noted that such searches may have been insensitive to FCHAMPs whose charge is close to an integer ($|Q_{L} - n| \\leq 0.25$) and we identified some islands in the \\ql~-- \\ml~plane (see Figure~\\ref{islands}) where fractionally charged massive particles may not have been excluded.  FCHAMPs with these charges and masses would reside in the ISM of the Galaxy and may have been incorporated into the cosmic rays.  Searches for such particles and, if there is evidence for them, the ratios of different charge states ($\\pm Q_{L}$, $+(2 - Q_{L})$, $+(1 - Q_{L})$, etc.) can provide a smoking gun for the existence of FCHAMPs.  It is also anticipated that the collider limits, especially for large \\ql\\ and \\ml, will be considerably strengthened at the LHC.  \\begin"
    },
    "1107/1107.1842_arXiv.txt": {
        "abstract": "{Investigations of mass segregation are of vital interest for the understanding of the formation and dynamical evolution of stellar systems on a wide range of spatial scales. A consistent analysis requires a robust measure among different objects and well-defined comparison with theoretical expectations. Various methods have been used for this purpose but usually with limited significance, quantifiability, and application to both simulations and observations.} {We aim at developing a measure of mass segregation with as few parameters as possible, robustness against peculiar configurations, independence of mass determination, simple implementation, stable algorithm, and that is equally well adoptable for data from either simulations or observations.} {Our method is based on the minimum spanning tree (MST) that serves as a geometry-independent measure of concentration. Compared to previous such approaches we obtain a significant refinement by using the geometrical mean as an intermediate-pass.} {The geometrical mean boosts the sensitivity compared to previous applications of the MST. It thus allows the detection of mass segregation with much higher confidence and for much lower degrees of mass segregation than other approaches. The method shows in particular very clear signatures even when applied to small subsets of the entire population. We confirm with high significance strong mass segregation of the five most massive stars in the Orion Nebula Cluster (ONC).} {Our method is the most sensitive general measure of mass segregation so far and provides robust results for both data from simulations and observations. As such it is ideally suited for tracking mass segregation in young star clusters and to investigate the long standing paradigm of primordial mass segregation by comparison of simulations and observations.} ",
        "introduction": "\\label{sec:introduction}  It is commonly accepted that star formation does usually not occur in isolation but that a large majority of young stars -- up to 90\\,\\% -- are part of a cluster \\citep{2003ARA&A..41...57L,2009ApJS..181..321E}. The dynamical evolution of a star cluster leaves a variety of imprints in the phase space of its stellar population which are good tracers of the \\emph{dynamical} age of the cluster. This quantity is in particular interesting when compared to the \\emph{physical} age. A higher dynamical than physical age means that observable dynamical imprints did not have enough time to evolve dynamically and thus must have been present - at least partially - already at the beginning. This is usually known as \\emph{primordial} origin.  One of the most widely discussed aspects of the dynamical evolution of young star clusters is that of mass segregation. From theoretical work it is well known that this process is inevitably entangled with the dynamical evolution of a self-gravitating system of at least two different mass components \\citep{1969ApJ...158L.139S,1983ApJ...271...11F,1995MNRAS.272..772S,2007MNRAS.374..703K}. Due to energy equipartition -- hence via two-body encounters -- the more massive particles tend to settle towards the cluster centre over time while the lower-mass particles are preferentially pushed to the outer parts. However, it is a much more challenging task to identify mass segregation observationally in real objects than theoretically from `clean' numerical simulations.  This is even more severe for young star clusters that are usually still embedded in their natal gas and the dynamical and physical age of which is much more difficult to estimate.  However, the investigation of mass segregation in these young stellar systems is of particular interest for a deeper understanding of the star formation process. Three fundamental questions are part of the scientific discussion in this context: (i) Do young star clusters (really) show mass segregation? (ii) What is the observed degree of mass segregation? (iii) Could the observed degree of mass segregation have developed dynamically or can it be only explained by primordial origin?  An investigation of these important aspects of the star formation process requires a solid tool that is at best independent of the method used to determine the stellar masses and independent of the geometry of the object, that provides an unambiguous measure and is equally well applicable to observational and numerical data. Note that also dynamical models have the equivalent problem to identify mass segregation clearly and quantitatively.  So far mainly four measures have been used to investigate mass segregation in (young) star clusters: i) the slope of the (differential) mass function in different annuli around the cluster centre \\citep[$\\Mmfd$; e.g.][]{1988ApJ...329..187R,1989ApJ...341..168B,1997AJ....113.1733H}, ii) the slope of the cumulative mass function in different annuli around the cluster centre \\citep[$\\Mmfc$; e.g.][]{1992BASI...20..287P,1998ApJ...492..540H}, iii) the characteristic radius of different mass-groups of stars \\citep[$\\Mrch$; e.g.][]{1983ApJ...271...11F}, and iv) the length of the minimum spanning tree (MST) of different mass-groups \\citep[$\\Mmstl$; developed by][]{2009MNRAS.395.1449A}\\footnote{A variant of the MST method has been presented just recently by \\citet{2011arXiv1103.0406Y} at the time of submission of our publication.}.  Most of these methods suffer from several weaknesses. The first three, $\\Mmfd$, $\\Mmfc$, and $\\Mrch$, implicitly assume a spherical geometry and thus depend on the determination of some cluster centre. The first two of them introduce additional bias due to radial binning and uncertainties in deriving the slope of the mass function. Furthermore, $\\Mmfd$ suffers from uncertainties due to mass binning \\citep[see e.g.][for a comparison of $\\Mmfd$ and $\\Mmfc$ applied to observational data]{2006AJ....132..253S}. There is a fundamental difference in the concept of the first and the last two methods: the former measure the mass distribution in different spatial volumes, the latter evaluate the spatial distribution of different sets of most massive stars. Consequently, $\\Mrch$ and $\\Mmstl$ do \\emph{not} require a direct measure of stellar masses but only a qualitative criterion for correct ordering. This property is a huge advantage in the face of observational data. However, $\\Mrch$ has also a significant weakness: the characteristic radius of a small subgroup is very sensitive to the definition of the cluster centre. Hence this method is not viable for a low degree of mass segregation, i.e. a signature from only a few most massive stars.  In contrast, $\\Mmstl$ does not show any of these disadvantages. However, it is also affected by quite low sensitivity that prevents unambiguous detection of weak mass segregation. To take advantage of the potential strength of $\\Mmstl$ we have developed an improved prescription for measuring mass segregation, in the following referred to as $\\Mmstg$, with significantly increased sensitivity. We present in this work our scheme and demonstrate its efficiency.  In Section~\\ref{sec:methods} we describe our method $\\Mmstg$ for measuring mass segregation and discuss a scheme for setting up initially mass segregated star cluster models. In Section~\\ref{sec:results} we discuss various examples of the effectiveness of our scheme compared to previous methods and present a first numerical application. The conclusion and discussion mark the last section of this paper. ",
        "conclusions": "\\label{sec:conclusions}  We have developed a new method, $\\Mmstg$, to measure mass segregation by significantly improving a previous approach of \\citet[][]{2009ApJ...700L..99A} based on the minimum spanning tree (MST), here referred to as $\\Mmstl$. Their method uses the normalised length of the MST of a given sample of stars, $\\Lmst$, as a measure of compactness.  Compared to ``classical'' measures of mass segregation such as the slope of the (differential) mass function in different annuli around the cluster centre (see Section~\\ref{sec:introduction}), $\\Lmst$ overcomes several substantial weaknesses. Our new method $\\Mmstg$ inherits all the advantages provided by $\\Mmstl$ -- \\begin{itemize} \\item independence of cluster geometry, \\item no requirement of cluster centre, \\item no requirement of quantitative measure of masses, \\item one-parametric unique measure of mass segregation -- \\end{itemize} and adds three significant improvements: \\begin{itemize} \\item nearly linear computational efficiency, \\item robustness against peculiar configurations, and \\item high sensitivity. \\end{itemize}  This gain is based on two boosting ingredients: \\begin{enumerate} \\item The implementation of an efficient ${\\cal{O}}(N \\log{N}$) algorithm for the calculation of the MST, and \\item the calculation of the geometrical mean of the connecting edges, $\\Gmst$, instead of just their sum, $\\Lmst$. \\end{enumerate}  The high performance is achieved by construction of a graph via 2D Delaunay triangulation \\citep[from the software package GEOMPACK:][]{1991AdvEng.13..325J}, quick sorting of its edges, and an efficient numerical implementation of Kruskal's algorithm \\citep{1956PAMS....7...48K}. The advantage of the geometrical mean is that it damps the contribution of ``outliers'' very efficiently and thus traces the concentration of the dominant stars of a configuration.  We have demonstrated the advantage of $\\Mmstg$ compared to other known and frequently used methods via various examples, both for numerical and observational data. \\begin{enumerate} \\item In general, using only the ten to twenty most massive stars $\\Gmst$ provides a robust and sensitive measure of mass segregation for the entire population of star clusters in our Galaxy. \\item In particular, very low degrees of mass segregation can be detected in massive clusters like NGC~3603 that consist of 10k or more stars. \\item We have also confirmed the very strong mass segregation of the five most massive stars in the ONC as reported by \\citet{2006ApJ...644..355H}. \\item Incompleteness of observational data has only a moderate effect on $\\Mmstg$ and corrections can be implemented easily similar to methods like $\\Mmfd$. \\end{enumerate}   However, we also stress three important aspects that have to be considered for a careful investigation of mass segregation via $\\Mmstg$: \\begin{enumerate} \\item Our conclusion that the 10 to 20 most massive stars generally provide the most sensitive measure of mass segregation has been derived from model clusters with smooth mass segregation over the entire sample. We caution that this is not to be expected if only specific subsamples are affected by mass segregation (as shown for the five most massive stars of the ONC in Section~\\ref{sec:results:onc}); this is of particular concern for studies of primordial mass segregation. We thus generally recommend to use both the cumulative and the differential analysis of a set of mass groups to investigate mass segregation in (young) star clusters. \\item The signature of a small sub-sample can be strong enough to bias $\\Gmst$ of a larger parent sample. To avoid wrong conclusions one should also calculate $\\Gmst$ of disjoint samples, e.g. of the 5, 6 to 10, 11 to 20 etc. most massive stars. \\item The value of $\\Gmst$ of a sub-sample is independent of its location with respect to the entire sample. It only reflects the compactness of this sub-sample, \\emph{not} its concentration towards some centre of the entire sample. So a large $\\Gmst$ does not necessarily mean that the corresponding sub-sample is mass-segregated.  Thus, in practice, in particular when dealing with observational data, the determination of some centre of the sample as reference point is still required. However, the value of $\\Gmst$ remains \\emph{independent} of the reference point and so its advantage compared to other methods. \\end{enumerate}  Our improved method $\\Mmstg$ will be vital for future studies of mass segregation in stellar systems. Its robustness and sensitivity is ideal for tracing lower degrees and more complex types of mass segregation than before. This is of particular interest for the earliest stages of star formation that seem to form protoclusters of complex structure \\citep[e.g.][]{2006ApJ...636L..45T}. Furthermore, it will also help to investigate the question whether mass segregation is observed at all in young star clusters \\citep[e.g.][]{2009A&A...495..147A}. Finally, our method can be used in general for precise structure analysis of any type of stellar system from planetary debris disks to galaxy clusters.  As a first application of $\\Mmstg$ to numerical simulations we have demonstrated that mass segregation in young star clusters can occur very quickly under dynamically cold initial conditions. Indeed we do expect young stellar systems to form under such conditions as a result of the gravitational collapse of their parent molecular cloud. Whether collapse still occurs in largely evolved young clusters like the ONC -- the scenario corresponding to our simulation -- is still unknown yet became favoured recently \\citep[e.g.][]{1988AJ.....95.1755J,2000NewA....4..615K,2005MNRAS.358..742S,2008ApJ...676.1109F,2009ApJ...697.1020P,2009ApJ...697.1103T}. However, at least some fraction of the oldest stars must have formed during global collapse. Hence rapid mass segregation affects the structure of star forming regions at the earliest stages and could thus explain the observed mass segregation of the most massive stars in the ONC.  We will investigate this and other aspects of dynamical mass segregation in young star clusters under various conditions in an upcoming publication."
    },
    "1107/1107.4105.txt": {
        "abstract": "Using \\emph{HST} observations of 147 host galaxies of low-mass black holes (BHs), we systematically study the structures and scaling relations of these active galaxies. Our sample is selected to have central BHs with virial masses $\\sim10^5-10^6\\msun$. The host galaxies have total $I$-band magnitudes of $-23.2<M_I<-18.8$ mag and bulge magnitudes of $-22.9<M_I<-16.1$ mag. Detailed bulge-disk-bar decompositions with GALFIT show that $93\\%$ of the galaxies have extended disks, $39\\%$ have bars and $5\\%$ have no bulges at all at the limits of our observations.  Based on the S\\'ersic index and bulge-to-total ratio, we conclude that the majority of the galaxies with disks are likely to contain pseudobulges and very few of these low-mass BHs live in classical bulges.  The fundamental plane of our sample is offset from classical bulges and ellipticals in a way that is consistent with the scaling relations of pseudobulges.  The sample has smaller velocity dispersion at fixed luminosity in the Faber-Jackson plane, compared with classical bulges and elliptical galaxies. The galaxies without disks are structurally more similar to spheroidals than to classical bulges according to their positions in the fundamental plane, especially the Faber-Jackson projection.  Overall, we suggest that BHs with mass $\\lesssim 10^6 \\msun$ live in galaxies that have evolved secularly over the majority of their history.  A classical bulge is not a prerequisite to host a black hole. ",
        "introduction": "\\label{sec:intro}  In the past decade, we have found strong correlations between supermassive black hole (BH) masses ($\\sim 10^6-10^9\\msun$) and the properties of host bulges (e.g, the $M_{\\rm BH}-\\sigma_{\\ast}$ relation; \\citealt{Tremaineetal2002}, and the $M_{\\rm BH}-L_{\\rm bulge}$ relation; \\citealt{MarconiHunt2003}). However, very little is known about BH-bulge correlations for low-mass BHs ($<10^6\\msun$) or late-type galaxies. The correlations between BH mass and galaxy properties in this low-mass regime provide important constraints on the formation mechanisms of the first primordial seed BHs \\citep[e.g.,][]{volonterinatarajan2009}.  Furthermore, low-mass BHs are expected to be a major source of gravitational radiation \\citep[e.g.,][]{hughes2002}.  Given the strong correlations between bulge properties and BH mass, the question remains whether supermassive BHs can exist in bulgeless galaxies, and if so how massive they become.  It was originally thought that BH mass was linked exclusively with bulge mass. For instance, the nearby late-type spiral galaxy M33 does not contain a central BH more massive than 1500 $\\msun$ \\citep[e.g.,][]{Gebhardtetal2001,Merrittetal2001}. Similarly, a central BH in the nearby dwarf galaxy NGC 205 has an upper limit of $2.2\\times10^4$ $\\msun$ on its mass \\citep[e.g.,][]{Vallurietal2005}. On the other hand, recent observations show that central massive BHs can also exist without classical bulges \\citep[e.g.,][]{FilippenkoHo2003,GreeneHo2004,Barthetal2004, Greenetal2007,Satyapaletal2007,Shieldsetal2008,Satyapaletal2009, Barthetal2009}.  Prior to our work, NGC 4395 and POX 52 were the only galaxies known to host BHs with $\\mbh\\lesssim 10^6~\\msun$.  NGC 4395 is an Sdm spiral with no bulge.  POX 52 is a spheroidal galaxy, also sometimes called a dwarf elliptical, although we follow the naming convention of \\citet{Kormendyetal2009}.  One has a disk, and the other has no disk.  Neither has a classical bulge component.  A larger sample of low-mass BHs is needed to understand the properties of their host galaxies.  Finding low-mass BHs is particularly challenging.  Unlike massive BHs in nearby galaxies, it is nearly impossible to measure dynamical BH masses for BHs with $\\lesssim10^6\\msun$ outside of the Local Group. We cannot yet resolve their gravitational spheres of influence, although it is possible to place interesting constraints on central BH masses in nearby objects \\citep[e.g.,][]{Barthetal2009,Sethetal2010}. Instead, we rely on indirect methods to estimate the ``virial'' masses of actively accreting BHs.  Based on the broad H$\\alpha$ profile and the calibrated radius-luminosity relation \\citep[e.g.,][]{Bentzetal2009}, \\cite{GreeneHo2004} presented 19 galaxies with virial BH masses $\\lesssim10^6\\msun$.  The sample has since increased to 174 galaxies (\\citealt{Greenetal2007}, see also \\citealt{Dongetal2007}).  From the SDSS data alone, we do not learn much about the host galaxy properties.  Here we present a study of the host galaxies of this large sample of low-mass BHs using \\emph{Hubble Space Telescope} (\\emph{HST}) observations.  Our primary goal is to determine the morphological types of galaxies hosting low-mass BHs, in particular the types of bulges that they contain. From the SDSS images, we know that the galaxies are $\\sim 1$ mag below $L^*$ \\citep{GreeneHo2004}.  In this luminosity range, galaxies have many different morphologies, ranging from small elliptical galaxies (e.g., M32) to late-type spirals and spheroidals. We want to determine whether a bulge is a necessary requirement to host a central supermassive BH.  First, we will measure the fraction of low-mass BHs without a bulge-like component of any kind. Bulgeless galaxies observed in the nearby universe, such as NGC 4395 and NGC 6946 (e.g., \\citealt{FilippenkoHo2003,Shihetal2003,Boomsmaetal2008}), are recognized as a challenge to the cold dark matter galaxy formation scenario \\citep[e.g.,][]{KormendyFisher2008,Kormendyetal2010,PeeblesNusser2010}. The fact that there are BHs in bulgeless galaxies implies that a bulge is not a necessary condition for the formation of BHs. \\cite{Satyapaletal2009} study $18$ truly bulgeless Sd/Sdm galaxies with {\\it Spitzer} and find only one active galaxy (NGC 4178), which suggests that BHs in bulgeless galaxies may be truly rare.  X-ray observations find active galactic nuclei (AGNs) in $\\sim 25\\%$ of Scd---Sm galaxies (\\citealt{DesrochesHo2009}). Larger samples of such bulgeless host galaxies will elucidate their nature.  Second, we will determine what fraction of the low-mass BH hosts have no disk.  These may be either elliptical or spheroidal galaxies \\citep[e.g.,][]{Ferrareseetal2006,Kormendyetal2009}.  More massive BHs are found in elliptical galaxies, but these lower-mass systems have stellar masses that are consistent with being either small ellipticals or spheroidals.  We will use their scaling relations (e.g., the fundamental plane) to distinguish between these two types.  The differing structures of spheroidal and elliptical galaxies likely reflects different formation histories \\citep[e.g.,][]{Kormendyetal2009}, and so the question is whether BH formation and growth occurs in spheroidal systems.  Third, for the disk galaxies with a bulge component, we ask whether they are classical bulges or not.  Observational evidence is building that there are two different kinds of bulges, namely classical bulges and pseudobulges (e.g., \\citealt{Kormendyetal2004}).  Pseudobulges have properties, including rotational support, exponential profiles, and ongoing star formation, that implicate the importance of secular processes such as bar transport in their build-up. We will measure the structure of the bulges from the \\emph{HST} observations to decide whether they are pseudobulges or classical bulges.  Finally, many recent papers have hinted that the scaling relations between low-mass BHs and their bulges are systematically different compared with more massive BHs \\citep[e.g.,][]{Hu2008,Greenetal2008, GadottiKauffmann2009, Greenetal2010}.  Ultimately, we will use the structural measurements presented here to examine the $M_{\\rm BH}-L_{\\rm bulge}$ and $M_{\\rm BH}-M_{\\rm bulge}$ correlations for this sample \\citep{Jiangetal2011}.  In \\S\\ref{sec:data}, we describe the data and reduction process.  In \\S\\ref{sec:image}, we present the image decompositions.  We explain how the uncertainties and upper limits are estimated in \\S\\ref{sec:uncertainty}.  Morphological results are given in \\S\\ref{sec:morphology} and scaling relations are given in \\S\\ref{sec:pseudobulge}. Finally, in \\S\\ref{sec:summary}, we summarize the paper.  The following cosmological parameters have been adopted: $H_0=100h=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$ and $\\Omega_{\\Lambda}=0.75$ (\\citealt{Spergeletal2003}).  Galactic extinction is calculated based on the fitting formula given by \\cite{Cardellietal1989}. ",
        "conclusions": "\\label{sec:summary}  We have looked at the host galaxies of 147 active galaxies selected to have low-mass ($\\mbh \\lesssim 10^6 \\msun$) BHs. Using \\emph{HST}/WFPC2, we perform detailed two-dimensional bulge, disk and bar decompositions of the entire sample, to study the bulge morphologies and structures of this unique sample. We find that the sample is dominated by disk galaxies (only $7\\%$ have no disk) with small bulge components.  We return to the questions we raised in the introduction.  Most host galaxies of low-mass BHs have a bulge component.  Only seven objects ($5\\%$) are consistent with being bulgeless galaxies. These are only candidates for bulgeless active galaxies because of the fit uncertainties, so BHs without any bulge component are apparently rare. We do have to keep in mind that we are biased by optical selection, which limits us to AGNs with high Eddington ratios in relatively massive galaxies.  Multi-wavelength approaches are required to truly determine the space density of AGNs in bulgeless galaxies.   Turning to those galaxies with bulges, the only remaining question is whether they are classical or pseudobulges.  Based on the low S{\\'e}rsic indices, low bulge-to-total ratios, and the prevalence of bars, rings, and nuclear spirals, we argue that the disk sample is dominated by pseudobulge galaxies.  Consistent with this supposition, we find that the fundamental plane of these galaxies is different from both elliptical galaxies and spheroidal galaxies, but consistent with observations of pseudobulge fundamental plane scalings.  In particular, the galaxies are larger and have lower velocity dispersions at a given luminosity than elliptical galaxies.  After we account for their young stellar populations, the differences are even more significant.  Our sample is also found to have a flatter Faber-Jackson relation in the $L-\\sigma$ plane.  Of the 147 galaxies, we only find 13 galaxies with nearby companions that are candidates for ongoing mergers, which is consistent with the fact that pseudobulges evolve via secular processes.  In conclusion, the host galaxies of low-mass BHs are different from classical bulges, and have properties that are consistent with pseudobulges.  The 11 bulges without extended disks do not scale as elliptical galaxies either. Their fundamental plane scaling relations are systematically offset, especially in the Faber-Jackson plane.  We have shown definitively that a classical bulge is not a prerequisite to host a supermassive BH.  We can also determine whether galaxies selected by their nuclear activity differ in any way from field galaxies selected at the same luminosity or mass. We have already seen that our disk galaxies have comparable bar fractions and merger fractions as inactive galaxies.  We now look at the distribution of morphological types for a more complete sample of galaxies at similar luminosity.  We use the morphology-dependent luminosity function from \\cite{Nakamuraetal2003}. For galaxies with $r^{\\ast}$-band magnitudes $M_r\\approx-18$ mag, which corresponds to an $I$-band magnitude $M_I \\approx-19$ mag, the relative number of elliptical/S0, Sa/Sb spiral galaxies and Sc/Sd spiral galaxies is found to be $1.00:4.10:0.73$. Our sample has relatively more late-type spiral galaxies than field galaxies selected by SDSS (Figure \\ref{BTratio}).  The simplest explanation for this observed difference is that we have a bias towards spiral galaxies because they have the gas fuel needed to feed a high-luminosity AGN.  Among spirals, inactive galaxies at the luminosity of our sample are dominated by pseudobulges. \\cite{Weinzirletal2009} finds that $\\sim76\\%$ of nearby high-mass spiral galaxies have low $n\\le2$ bulges -- they are  pseudobulges. In a sample of 173 E-Sd galaxies, \\cite{FisherDrory2010} also finds that over $78\\%$ of the identified bulges are pseudobulges. This implies that our AGN-selected sample has very similar bulge properties to non-AGN selected samples.  Our results help elucidate the growth mechanisms of low-mass BHs, as they are found preferentially in pseudobulges, which are thought to be evolving secularly with a quiescent recent history \\citep[e.g.,][]{Kormendyetal2004}.  Most likely, these low-mass BHs are not fueled by major mergers, which is also consistent with the fact that only $9\\%$ of our galaxies have close companions.  Rather, the secular processes that build up the bulge may well fuel the AGN as well.  In fact, it is possible that these BHs were formed with a mass quite similar to their present mass (e.g., \\citealt{PortegiesMcmillan2002,Koushiappasetal2004,Begelmanetal2006}). If supermassive BHs indeed grow from low-mass BHs, some violent event (e.g., a merger) is likely required to both dramatically change the bulge properties and substantially grow the BH.  As already discussed in \\cite{Greenetal2007} and \\cite{Greenetal2010}, differences in the bulge properties of low-mass BHs lead to differences in the $M_{\\rm BH}$ --- $\\sigma_{\\ast}$ and $M_{\\rm BH}$ -- $L_{\\rm bulge}$ relations at low mass (if these relations still exist at all). The $M_{\\rm BH}$ -- $\\sigma_{\\ast}$ relation of this sample is discussed in more detail in \\cite{Xiaoetal2011} while the $M_{\\rm BH}$ --- $L_{\\rm bulge}$ relation is discussed in a companion paper \\citep{Jiangetal2011}. In that paper, we find that the tight scaling relations between BH mass and bulge luminosity found for massive classical bulges do not exist for the pseudobulges that we have shown are the hosts of low-mass BHs."
    },
    "1107/1107.0596_arXiv.txt": {
        "abstract": "{\\\\  Context: The origin of the variable component of the solar wind is of great intrinsic interest for heliophysics and space-weather, e.g. the initiation of coronal mass ejections, and the problem of mass loss of all stars. It is also related to the physics of coronal neutral sheets and streamers, occurring above lines of magnetic polarity reversal.  Filaments and prominences correspond to the cool coronal component of these regions.\\\\  Aims: We examine the dynamical behaviour of these structures where reconnection and dissipation of magnetic energy in the turbulent plasma are occurring. The link between the observed oscillatory motions and the eruption occurrence is investigated in detail for two different events.\\\\  Method: Two filaments are analysed using two different datasets: time series of spectra using a transition region line (He~I at 584.33 ${\\rm \\AA}$) and a coronal line (Mg~X at 609.79 ${\\rm \\AA}$) measured with CDS on-board SOHO, observed on May 30, 2003, and time series of intensity and velocity images from the NSO/Dunn Solar Telescope in the H$\\alpha$ line on September 18, 1994 for the other. The oscillatory content is investigated using Fourier transform and wavelet analysis and is compared to different models.\\\\  Results: In both filaments, oscillations are clearly observed, in intensity and velocity in the He~I and Mg~X lines, in velocity in H$\\alpha$, with similar periods from a few minutes up to 80 minutes, with a main range from 20 to 30 minutes, cotemporal with eruptions. Both filaments exhibit vertical oscillating motions. For the filament observed in the UV (He I and Mg X lines), we provide evidence of damped velocity oscillations, and for the filament observed in the visible (H$\\alpha$ line), we provide evidence that parts of the filament are oscillating, while the filament is moving over the solar surface, before its disappearance. ",
        "introduction": "Prominences (observed above the limb) and filaments (observed on the solar disk) are complex structures of cool, dense plasma embedded in the 2 orders of magnitude less dense corona, which is also 2 orders of magnitude hotter. They are maintained in dynamical equilibrium against gravity by the magnetic field, and they can channel waves propagating outward from the low chromosphere. Filaments and prominences form in both active and quiet-Sun regions. Filaments and prominences that have been stable for weeks can suddenly erupt (e.g. \\cite{Oliver_09}). Some of the erupting filaments, heated and becoming visible in coronal emission lines, do not produce a CME: the filament reforms after cooling due to radiative losses. Such an example was reported by \\citet{Filippov_02} who discuss the possible origin of the heating related to the dynamical behaviour of the filament, including oscillations.\\\\  As pointed out by \\citet{Isobe_Tripathi_06} and \\citet{Tripathi_09}, oscillations of erupting filaments provide an alternative tool for probing the onset mechanisms of eruptions. SOHO/SUMER observations of a prominence oscillation were discussed in the context of the physics of CMEs \\citep{Chen_08}. It was suggested that repetitive reconnections between emerging flux and the pre-existing magnetic field are the trigger mechanism of Doppler velocity oscillations leading to the eruption. Periods near 20 minutes, lasting 4 hours before the eruption, were detected. However, much of the literature deals only with quiescent filaments or prominences. The oscillations are classified in two groups based on their amplitude: large amplitudes of several tens of km/s  \\citep{Ramsey_66}, and small amplitudes of a few km/s or less \\citep[e.g.][]{Terradas_02, Oliver_02, Engvold_08}. Large amplitude oscillations in filaments, typically 20 km/s or more, are rare events \\citep{Tripathi_09}; they are due to flares or to a local emergence of magnetic flux near the filament, while small amplitude oscillations do not affect the whole structure at a time but are of a local nature; intermediate values are also reported \\citep[e.g.][]{Bocchialini_01} with velocities up to around 13 km/s.\\\\  The oscillations in quiescent prominences are classified into three categories traditionally measured from their Doppler signal: short-period (less than 5 minutes), intermediate periods (in the range of 6 to 20 minutes) detected in individual threads of prominences, and long-period (between 40 to 90 minutes) detected over the whole prominence or filament \\citep{Molowny_97, Oliver_02, Engvold_08, Mackay_10}. \\citet{Ning_09} report threads in prominences exhibiting vertical and horizontal motions from Hinode/SOT observations close to the H$\\alpha$ line center at 6563+0.076 ${\\rm \\AA}$; the predominant period of the small-amplitude velocity oscillations is about 5 minutes. \\citet{Berger_10} recently reported small-scale turbulent upflows in quiescent prominences observed with Hinode/SOT, in the Ca II H line at 3969 ${\\rm \\AA}$ and the H$\\alpha$ line at 6563 ${\\rm \\AA}$. The recurrence time of these upflows is roughly 5-8 minutes and the location remains active on timescales from tens of minutes up to several hours. \\citet{Isobe_Tripathi_06} and \\citet{Isobe_07} report the oscillatory motion of a polar crown filament, observed at 195 ${\\rm {\\rm \\AA}}$ at a 12 minute cadence and in H$\\alpha$, with a long and constant period of about 2 hours: 5 km/s amplitude oscillations, measured in the plane of the sky, are detected in 195-${\\rm {\\rm \\AA}}$ images, and occured during the slow rise of the filament, which erupted after 3 periods. \\citet{Pinter_08} observed the same filament using the same EUV dataset and wavelet spectra to extract intensities, which show that the filament oscillates as a rigid body, with a period of about 2.5-2.6 hours, which is constant along the filament. All the above analysis do not reveal any signature of the eruption. Recent detection of long and ultra-long period oscillations of up to several hours have been reported: \\citet{Foullon_04} report the first detection of 8-27 hour period in oscillatory intensity variations of a quiescent filament observed at 195 ${\\rm {\\rm \\AA}}$. The dominant period is 12.1 hours. \\citet{Pouget_06} report 5-6 hour period oscillations detected in the Doppler velocity of a quiescent filament, in He~I at 584.33 ${\\rm {\\rm \\AA}}$. \\citet{Foullon_09} show that the long period detected in a quiescent filament, using 195 ${\\rm {\\rm \\AA}}$ images, increases prior to its eruption.\\\\  Small-amplitude oscillations have been extensively modelled over the past 20 years in parallel with the development of coronal seismology \\citep{Ballester_06}, which aims to interpret the oscillations detected in the corona and its structures in terms of MHD waves as a tool to diagnose those structures. However, most of these MHD models of prominence oscillations deal with stable filaments and cannot predict their eruption. \\citet{Vrsnak_90a} developed a simple model of cylindrical prominences to study their eruptive instability and provided a scenario for the eruption onset in terms of mass loss, increase of the electric current caused by slow reconnection below the filament, and twisting motions in the footpoints. Indeed, \\citet{Koutchmy_08} observed such behaviour in a prominence just before it erupted and gave rise to a CME, compatible with the standard model \\citep{Priest_00}, which invoke resistive instabilities. \\citet{Filippov_08} proposed a scenario where the slow rising of a filament is interpreted as an increase of the non-potentiality of the prominence magnetic field inside the filament before its eruption, when compared to the potential field in the surrounding corona, computed from the observed surface fields.\\\\  Several authors \\citep[e.g.][]{Wiehr_89, Molowny_98,Terradas_02, Vrsnak_07, Ning_09} discussed damped oscillations in quiescent filaments, disappearing after a few periods, i.e. quickly damped by one or several mechanisms \\citep{Oliver_09}. \\citet{Arregui_10} reviewed the theoretical damping mechanisms in quiescent prominences as thermal effects, through non-adiabatic processes, mass flows, resonant damping and partial ionization effects. \\citet{Vrsnak_90b} reported damped post-eruptive oscillations in an active prominence, however clear observations of prominences at limb during their eruption are extremely rare \\citep{Berger_10}. We focus here on the analysis of two active filaments, observed during their eruption, from two sets of observations performed at high cadence: one with the Coronal Diagnostic Spectrometer (CDS) \\citep{CDS_95} on-board SOHO \\citep{Fleck_95} and one with the NSO/Dunn Solar Telescope (DST). These data sets are complemented by observations from the extreme UV imager EIT on-board SOHO and the soft X-ray telescope (SXT) on-board Yohkoh. The two data sets are thus quite different: one consists of spectroscopic observations in the UV (SOHO/CDS), the other of narrow band filtergrams taken in the H$\\alpha$ line (DST).  However in both cases, wavelet analysis of the observed signals in the two filaments leads to evidence of intermediate-period, small-amplitude oscillations (less than 20 km/s) apparently cotemporal with a total or partial eruption of the filaments. Moreover, the CDS observations in He~I at 584.33 ${\\rm \\AA}$ and in Mg~X at 609.79 ${\\rm \\AA}$ reveal intensity oscillations in phase in both lines, preceded by velocity oscillations in phase in both lines. Both data sets reveal a vertical oscillatory motion inside the filaments: the CDS observations in He~I at 584.33 ${\\rm \\AA}$ exhibit a damped oscillation. The DST 2D imaging observations in the H$\\alpha$ wings reveal that parts of the filament are moving vertically in opposite directions, while the filament is moving over the solar surface, before its disappearance; no damping is measured for this event. In addition, a flare and a loop eruption are observed in Yohkoh/SXT. ",
        "conclusions": "We presented the results of the analysis of two eruptive filaments exhibiting some similarities observed in completely different manners: one observed in He~I and Mg~X with a spaceborne spectrometer, the other observed in H$\\alpha$ wings with a groundbased imager. In both cases, the eruptive filaments show small amplitude oscillatory motions associated with an eruption and a partial disappearance: the filaments lost a part of their mass, are heated and evaporate.  In He~I, the pseudo-period of the velocity oscillations is increasing from 20 to 50 minutes while the amplitude is decreasing: this damped velocity oscillation could be interpreted, as suggested by \\cite{Vrsnak_90b}, as the signature of the eruption, due to the weakening of the restoring force giving rise to the increasing period. In this model dealing with the oscillatory motions in an eruptive filament, the authors explain that once the filament is in its eruptive phase, the filament jumps from an unstable position to a stable one at a higher altitude. In this stable position, if the damping factor (inverse of the damping time) is less than the frequency of the first oscillation mode, the filament will oscillate around its equilibrium position, as a damped oscillator. This behaviour corresponds to that described by our observations in Section 2. An increased heating of the top part of the filament is another possibility suggested by \\cite{Filippov_02}.\\\\ We can link this increase of the period of the oscillating filament to the decreasing stiffness of an oscillating spring; this suggesting the weakening of the restoring force, which may also be related to the destabilisation and onset of the filament eruption which has modified the whole structure of the filament and thus the restoring forces \\citep{Isobe_Tripathi_06}.\\\\ The value of $\\tau/T$ = 1.9 obtained in our study may be compared to that obtained for a similar behaviour in a longitudinally damped oscillation of a filament observed in H$\\alpha$ by \\citet{Vrsnak_07}: $\\tau/T$= 2.3. They interpreted their observations in terms of poloidal magnetic flux injection by magnetic reconnection at one of the filament legs. They also pointed out an increase with time of the oscillation period and interpreted this behaviour in terms of nonlinear effects in the framework of a harmonic oscillator model.\\\\  Results based on the velocity in He~I and Mg~X indicate that the corona in which the filament channel is embedded responds coherently with the oscillating filament. The intensity signals in both lines are very closely in phase, showing nearly simultaneous variations. Moreover, the peak of intensity seen in both lines appears after the beginning of the oscillations detected in velocity, which indicates that the dynamical effects precede the thermal effects.\\\\  For the second case, the analysis of the velocity in H$\\alpha$, shows a large power peak detected for zone 1 in the extreme eastern part of the filament rooted in the chromosphere: the small values of the period are close to the typical periods of the photospheric (5 minute) and chromospheric (3 minute) oscillations. These oscillations are detected before the motion of the filament toward the Northwest and before its eruption. We suggest that the filament is excited by the photospheric and chromospheric oscillations \\citep{Bocchialini_01}, which are at the origin of the motion of the filament, with a triggering of the eruption and the flare due to reconnections. This result can be compared to the 4 minute periods obtained in H$\\alpha$ \\citep{Schmieder_91} in chromospheric filamentary structures, interpreted either in terms of eigenmodes excited by the chromospheric oscillations, or in terms of Alfv\\`en waves generated by the chromosphere and transmitted to the filamentary structures, as confirmed by the Hinode/SOT observations \\citep{Okamoto_07}.\\\\  Indeed, these proper motions are difficult to interpret as definite horizontal velocities. It seems rather that we are dealing with increasing turbulent proper motions with increasing ionisation ratio (temperature) leading to the disappearance of a large part of the filament after $t$= 54 minutes. The analysis of the velocity, for both datasets, revealed the presence of oscillations with a period from 20 to 30 minutes in He~I, Mg~X, and H$\\alpha$, cotemporal with the eruption of the filaments. These periods can be compared to the 20 minute period found by \\citet{Anzer_09} who discussed the different oscillation periods for global modes of quiescent prominences. Anzer used simple 1D prominence configurations to describe the magnetohydrostatic equilibrium and their small amplitude oscillations. The restoring force along the $z$-direction, perpendicular to the solar surface, resulting from the magnetic tension of the stretched field lines, leads to periods of 20 minutes in this direction. Our results can also be compared to those of \\citet{Terradas_05} who showed that for a typical prominence temperature of 6000\\,K, the oscillation period computed for a quiescent prominence is around 25 minutes and the damping time is around 19 minutes. The damping time increases with temperature and depends on the radiation time. \\citet{Ballester_06} points out in his review on solar corona seismology that other mechanisms are possible such as wave leakage, resonant absorption and ion-neutral collisions. \\citet{Soler_07} and \\citet{Soler_09} demonstrate that radiative effects of the prominence plasma are responsible for the attenuation of slow modes, connected to intermediate period (10 to 40 minutes) prominence oscillations. The simultaneous upflows and downflows detected in the filament observed in the wings of H$\\alpha$ could be compared to the results obtained by \\citet{Berger_10}: high spatial resolution observations, in H$\\alpha$ with Hinode/SOT, of quiescent prominences reveal the presence of upflowing plumes. Their occurrrence combined with the plume dimensions indicate that ``the plumes are sources of upward mass flux into the prominence which partially offset the continual downward mass flux of heavy prominence material''.\\\\  These results on small amplitude oscillations of the intensity and of the Doppler velocity signals in eruptive filaments contribute to the statistical study needed to check the link between oscillations and eruption precursors."
    },
    "1107/1107.2593_arXiv.txt": {
        "abstract": "\\noindent   I briefly review the physics of  cosmic bubble collisions in false-vacuum eternal inflation.  My purpose is to provide an introduction to the subject for readers unfamiliar with it, focussing on recent work related to the prospects for observing the effects of bubble collisions in cosmology.  I will attempt to explain the essential physical points as simply and concisely as possible, leaving most technical details to the references.  I make no attempt to be comprehensive or complete.  I also present a new solution to Einstein's equations that represents a bubble universe after a collision, containing vacuum energy and ingoing null radiation with an arbitrary density profile. ",
        "introduction": "This work reviews the physics of cosmic bubble formation and collisions, with a focus on recent work. It will be as self-contained as possible while avoiding technical details, with references to the literature where further details can be found.  I will use natural units ($8 \\pi G = c = \\hbar = 1$) throughout.  Another review of cosmic bubble collisions is \\cite{ajreview}.   \\subsection{Overview}  First-order phase transitions are ubiquitous in physics.  During a first-order transition, a meta-stable phase (the ``false'' vacuum) decays to a lower energy phase (the ``true'' vacuum, which may itself be either metastable or truly stable) either by quantum tunneling or because the decay is stimulated by some external influence.  The transition begins in a finite region---a bubble.  If the bubble is sufficiently large when it forms, it  expands into the false vacuum and collides with other bubbles.  Therefore in a static spacetime, the transition will eventually percolate and the false vacuum will disappear entirely.  The situation is different when gravity is included.  The energy density in a meta-stable phase does not change with the expansion of space; it is a cosmological ``constant''.  Therefore, if the false vacuum has positive energy it will undergo cosmic inflation---the volume of space containing it will expand exponentially with time.  If the timescale for the exponential (the false-vacuum Hubble rate) is faster than the rate of bubble formation, the false vacuum will continually reproduce itself and the transition will never percolate.  Bubbles of true vacuum will expand (but slower than the false vacuum) and occasionally collide with each other (Fig. \\ref{fractal}).  Contained inside each bubble is an expanding Friedmann-Robertson-Walker (FRW) cosmology that is  homogeneous and isotropic apart from random perturbations and the aftereffects of collisions.  In  models consistent with current observations, our observable universe is  inside such a bubble, embedded in and expanding into an eternally inflating parent false vacuum.  Observational constraints require that our bubble underwent its own period of slow-roll inflation after its formation.  Collisions with other bubbles that nucleate in the (otherwise eternal) parent false vacuum nearby occur with a non-zero probability per unit time, and are therefore guaranteed to happen eventually.  Their effects have already or will perturb the universe around us, creating potentially observable signals.  In terms of the FRW coordinates describing the cosmology inside our bubble these collisions occur {\\it before} slow-roll inflation---in fact, they occur before the (apparent) big bang of our FRW universe.  They can be regarded as creating a special and predictable set of initial conditions at FRW time $t=0$.  \\begin{figure} \\hspace{0 in}\\includegraphics[width=1\\textwidth]{fractal.png} \\caption{\\label{fractal} A numerical simulation showing the spatial distribution of bubbles at a late time in an eternally inflating false vacuum.  Bubbles that appeared earlier expanded for longer and are larger, but the physical volume of the false vacuum is larger at later times, so there are more smaller bubbles.  Eventually, each bubble collides with an infinite number of others.  The late-time distribution is a scale-invariant fractal.} \\end{figure}   \\subsection{Effective field theory coupled to gravity}  Any model of spacetime fields coupled to gravity can give rise to bubble collisions if there exist at least two meta-stable phases of the field theory.   After forming, bubbles expand and  collide with each other.  The goal of this review is to describe the physics of the formation, expansion, and collision of these bubbles.  I will focus exclusively on models in which at least one of the phases (the false vacuum) has a positive vacuum energy.  A region of spacetime filled with positive vacuum energy has a metric \\be \\label{frwmet} ds^2 = -dt^2 + a(t)^2 d\\vec x^2 \\ee and obeys Einstein's equations \\be \\label{FRW} (\\dot a/a)^2=H_f^2 = {V_f / 3}, \\ee where $V_f$ is the energy density of the false vacuum and $H_f$ is the associated Hubble constant.  The solution to \\eqref{FRW} is de Sitter space, $a(t)=e^{H_f t}$.  Regions of space that are dominated by vacuum energy but contaminated by other forms of matter or energy will exponentially rapidly inflate away the contaminants and approach the metric \\eqref{frwmet}.  Regions not dominated by vacuum energy will either expand more slowly or collapse into black holes, which for many purposes is taken as justification for ignoring them after a few false vacuum Hubble times (where the Hubble time is $t_{H} \\equiv 1/H$).   \\subsection{Decay}  In metastable de Sitter space there is a dimensionless rate of bubble formation $\\gamma$.   When $\\gamma$ is small it can be defined as the expected number of bubble nucleations per unit Hubble time per unit Hubble volume:  that is, the dimensionful decay rate is $\\Gamma=H_f^4 \\gamma$.  Generally $\\gamma$ is the exponential of $-S$, where $S$ is the action for an instanton.  Hence, when $S \\gg 1$, $\\gamma$ is very small and the rate of bubble nucleations is  slow.  When $\\gamma \\simgeq 1$, the semi-classical methods  reviewed here are not adequate to describe the physics.  When $\\gamma$ is small and the meta-stable phase has positive vacuum energy, the exponential expansion means that the transition will never percolate---there will always be some regions in which the unstable phase remains.  The intuitive reason is simple:  in one Hubble time, the de Sitter region increases its volume by a factor of $e^{3}$.  In that same Hubble time, one expects $\\gamma$ bubbles of another vacuum, each of volume less than the Hubble volume, to appear.  Therefore if $\\gamma \\ll 1$ a typical region will double in size many times before producing a bubble.  The same applies to each ``child\" region; and therefore the meta-stable phase never ceases to exist.  This is known as ``false vacuum eternal inflation\" (to distinguish it from slow-roll eternal inflation, which can take place in an models with sufficiently flat, positive potential energy functionals).   In the end, the picture is an exponentially rapidly expanding spacetime in which bubbles of more slowly expanding phases occasionally appear, and occasionally collide.  It is important to note that this is a generic prediction of any model with multiple, positive energy minima---while it seems to be a prediction of the string theory landscape, it is certainly not unique to it.  Nevertheless, an observation that confirmed this model would be a confirmation of a prediction of string theory, and an observation that ruled it out would be at least potentially a falsification.   \\subsection{Motivation} \\label{mot}  In the last few years there has been a surge of interest in this problem.  The reason is that string theory predicts the existence of many meta-stable minima, the so-called ``string landscape\" \\cite{discretum, lennylandscape}.  In string theory, the geometry and topology of spacetime is dynamical.  String theories exist in 9 spatial dimensions.  Since we observe only three spatial dimensions, in string solutions that might describe our world six of the spatial dimensions are compactified; that is, they form a geometry with finite volume, while the other three spatial dimensions and time can form Minkowski or de Sitter space.  Six dimensional manifolds have many parameters that describe their shape and size.  These parameters are dynamical fields in string theory, and their meta-stable solutions correspond to distinct possibilities for the shape and size of the manifold.  Because they are meta-stable, small fluctuations around these geometries behave like massive particles from the point of view of the 4 large spacetime dimensions.  Therefore the low-energy physics as measured by a 4D observer is an effective field theory coupled to gravity, with the field content partially determined by which configuration the compact manifold is in.  The compact manifold can make transitions from one meta-stable configuration to another. Among the parameters that can vary from configuration to configuration is the value of the vacuum energy.  In non-supersymmetric solutions it will not be zero; instead, its typical value is set by the string energy scale---with both positive and negative values possible \\cite{discretum, kklt, lennylandscape}.   The exact value in any given phase will depend on the details, and if there are enough different phases, it is plausible that a small but non-empty subset have vacuum energies that are consistent with the tiny value we observe \\cite{riess, perlmutter, Komatsu:2010fb}.  Hence the theory is at least consistent with observation, and the extraordinarily small but non-zero value was even predicted a decade in advance in \\cite{w} under a set of assumptions that coincide with those just outlined.  Because understanding the value of the cosmological constant is considered one of the most important problems in theoretical physics, and because string theory is our best candidate for a theory of quantum gravity, it is worth taking this scenario seriously.  One should look for  observations beyond the cosmological constant which could test its validity; hence this review. ",
        "conclusions": "Very recently, a search of the WMAP temperature data \\cite{Feeney:2010dd, Feeney:2010jj} uncovered several cold and hot spots that are consistent with the signal predicted in \\cite{wwc2}.  While further analysis is necessary to determine whether or not these features are truly anomalous, this result raises the stakes significantly.  A discovery of a bubble collision in the CMB would confirm a prediction of string theory, and more broadly demonstrate the reality of eternal inflation, revolutionize our understanding of the big bang, and indicate the existence of other ``universes\".  With such possibilities in sight, it is well worth pursuing this program vigorously."
    },
    "1107/1107.5466_arXiv.txt": {
        "abstract": "Obscuration due to Galactic emission complicates the extraction of information from cosmological surveys, and requires some combination of the (typically imperfect) modeling and subtraction of foregrounds, or the removal of part of the sky. This particularly affects the extraction of information from the largest observable scales. Maximum-likelihood estimators for reconstructing the full-sky spherical harmonic coefficients from partial-sky maps have recently been shown to be susceptible to contamination from within the sky cut, arising due to the necessity to band-limit the data by smoothing prior to reconstruction. Using the WMAP 7-year data, we investigate modified implementations of such estimators which are robust to the leakage of contaminants from within masked regions. We provide a measure, based on the expected amplitude of residual foregrounds, for selecting the most appropriate estimator for the task at hand. We explain why the related quadratic maximum-likelihood estimator of the angular power spectrum does not suffer from smoothing-induced bias. ",
        "introduction": "It is unavoidable that we observe the Universe through the galaxy we inhabit. The foreground contamination injected by the Milky Way into full-sky cosmological data-sets must be modeled and removed, or the regions most conspicuously contaminated must be excised. Where no precise model of the foregrounds is available, cutting the sky is the most robust option, with the regrettable consequence that part of the signal is discarded along with the contamination. This includes information on the largest scales, which are valuable for a variety of reasons, including measurement of the integrated Sachs-Wolfe effect~\\cite{1967ApJ...147...73S} and constraining primordial non-Gaussianity using tracers of large-scale structure~\\cite{Slosar:2008hx}.  It is impossible to uniquely recover the cosmological signal discarded in the sky cut. However, by writing down the likelihood for the region of the sky in which one trusts the data, it is possible to reconstruct an estimate of the signal at large scales which maximizes the likelihood of the residual noise~\\cite{deOliveiraCosta:2006zj}. An alternative reconstruction scheme maximizes the {\\em posterior} probability~\\cite{Wiener:1964, Bunn:1994xn, Zaroubi:1994mx, Tegmark:1996qs, Bielewicz:2004en} of measuring the underlying cosmological signal given the available data  and a prior theoretical expectation on the signal.  The reconstructions estimate the large-scale (low-$\\ell$) spherical harmonic coefficients, $\\alm$, by treating the signal at small scales as noise and only considering data external to the sky cut. If, as with the cosmic microwave background (CMB), the field to be reconstructed is not band-limited, the proliferation of small-scale signal makes the reconstruction noisy to the point of being useless. Input maps are therefore smoothed -- necessarily prior to cutting the sky -- to truncate the signal and remove sources of confusion below a chosen angular scale~\\cite{Efstathiou:2009di}. However, smoothing leaks contamination from the masked region into the trusted data~\\cite{Aurich:2010gw,Copi:2011pe}, and the reconstructed spherical harmonic coefficients, $\\almest$, are biased. In this work we explore the causes and expected magnitudes of this bias, and discuss how it can be mitigated. ",
        "conclusions": "Maximum-likelihood estimators, $\\almest$, are often used to reconstruct the large-scale spherical harmonic coefficients, $\\alm$,  from partial-sky data. The technique relies on smoothing to restrict the amount of small-scale noise accessible to the reconstruction, but smoothing has been shown to contaminate ``clean'' pixels with residual foregrounds from within the sky cut. In this work, we have examined the impact of this smoothing-induced bias on the maximum-likelihood reconstruction. We have shown that it is possible to mitigate the bias by removing the contaminated regions, but these are only well-defined if smoothing is performed using a kernel with finite support on the sky. This precludes the use of the commonly used Gaussian kernel.  Cutting a larger portion of the sky greatly increases the variance of the reconstruction, but it is possible to counteract this effect by enforcing a prior on the reconstructed coefficients using a Wiener filter. We have therefore proposed an estimator -- using top-hat smoothing, extended masks and a Wiener-filtered reconstruction -- which does not suffer from smoothing-induced bias. By considering the expectation of the square of the reconstruction error, $\\ztot = \\sum_{\\ell,\\,m} \\langle (\\almest - \\alm)^2 \\rangle$, we have compared the performance of the maximum-likelihood and Wiener-filtered estimators in the presence of simulated CMB foreground residuals.  The reconstruction performance measure $\\ztot$ scales with the estimators' bias and variance, which in turn are governed by the amplitude of contamination and the size of the sky cut, respectively. The fiducial maximum-likelihood reconstruction is performed using relatively small sky cuts, but is susceptible to contamination through smoothing-induced bias; the finite-smoothing Wiener-filtered reconstruction does not suffer from smoothing-induced bias, but makes use of extended masks. Increasing the level of contamination therefore increases $\\ztot$ for the maximum-likelihood reconstruction only, which suggests that there is a level of contamination above which one should switch from the maximum-likelihood to the Wiener-filtered reconstruction.  Given an estimate of the morphology and amplitude of the contaminants within the cut sky, one can predict which modes will be biased and by how much, and hence determine the threshold at which one should swap estimators. We find that this threshold is relatively insensitive to the precise morphology of foreground residuals at large scales, and is mainly governed by their amplitude. Calculating $\\ztot$ for the two estimators in the presence of estimated foreground residuals, we determine this threshold to be $\\sim 10$ times the amplitude of the foreground residuals used in this work. Assuming that the ILC contains similar levels of contamination to those used here, we therefore recommend the use of either the contaminated full-sky $\\alm$s or the fiducial maximum-likelihood $\\almest$s when handling this data-set. However, when using foreground-reduced maps for individual WMAP frequencies, which contain much greater foreground residuals, the Wiener-filtered reconstruction will provide the best estimate of the large-scale underlying CMB signal. Note that, as the Wiener-filtered $\\almest$s are a maximum-{\\em posterior} solution, care must be taken if the reconstruction output is being used for further model-selection steps. The reconstruction techniques are, however, most commonly used to test the null hypothesis, in which case the prior employed in this work is completely appropriate.  For problems requiring only a power spectrum (as opposed to the full temperature field) the issues described in this paper are essentially irrelevant because the smoothing can be conducted on vastly smaller scales, the resulting range of poorly constrained modes being automatically downweighted."
    },
    "1107/1107.3183_arXiv.txt": {
        "abstract": "Results of a detailed abundance analysis of the solar twins \\cyga\\ and \\cygb\\ based on high-resolution, high signal-to-noise ratio echelle spectroscopy are presented.  \\cygb\\ is known to host a giant planet while no planets have yet been detected around \\cyga.  Stellar parameters are derived directly from our high-quality spectra, and the stars are found to be physically similar, with $\\Delta$\\teff\\ $=+43$ K, $\\Delta \\log g=-0.02$ dex, and $\\Delta \\xi=+0.10$ km s$^{-1}$ (in the sense of A $-$ B), consistent with previous findings. Abundances of 15 elements are derived and are found to be indistinguishable between the two stars.  The abundances of each element differ by $\\leq 0.026$ dex, and the mean difference is $+0.003 \\pm 0.015$ ($\\sigma$) dex.  Aside from Li, which has been previously shown to be depleted by a factor of at least 4.5 in \\cygb\\ relative to \\cyga, the two stars appear to be chemically identical.  The abundances of each star demonstrate a positive correlation with the condensation temperature of the elements (\\tc); the slopes of the trends are also indistinguishable.  In accordance with recent suggestions, the positive slopes of the [m/H]-\\tc\\ relations may imply that terrestrial planets have not formed around either \\cyga\\ or \\cygb.  The physical characteristics of the 16 Cyg system are discussed in terms of planet formation models, and plausible mechanisms that can account for the lack of detected planets around \\cyga, the disparate Li abundances of \\cygab, and the eccentricity of the planet \\cygb\\ b are suggested. ",
        "introduction": "\\label{s:intro} \\cyga\\ and \\cygb\\ are a well known common proper-motion pair of solar-twin stars with spectral types G1.5V and G3V, respectively.  Stellar parameters and [Fe/H] abundances of the pair have been derived by numerous groups \\citep[e.g.,][]{1994PASP..106.1248G,1996A&A...311..245F,1998A&A...336..942F,2001ApJ...553..405L,2005PASJ...57...83T}, and the abundances of additional elements have been derived by others \\citep[e.g.,][]{1993A&A...274..825F,1997AJ....113.1871K,1998A&AS..129..237F,1998A&A...334..221G,2000AJ....119.2437D,2001PASJ...53.1211T,2003MNRAS.340..304R,2004ARep...48..492G}. In each study, \\cygab\\ have been found to be physically similar, with A being slightly hotter and having a slightly lower surface gravity than B, consistent with their spectral types.  Differences in the derived stellar parameters in the sources listed above range from +25 to +62 K in \\teff, -0.03 to -0.15 dex in $\\log g$, and 0 to +0.05 dex in [Fe/H] (all comparisons herein are made in the sense of A $-$ B).  A defining property distinguishing the two stars is the designation of \\cygb\\ as a planet host.  \\citet{1997ApJ...483..457C} reported the presence of a radial-velocity detected planet (\\cygb\\ b) with $M \\sin i=1.5 \\; M_{\\mathrm{Jup}}$ orbiting \\cygb\\ on an eccentric orbit ($e=0.63$), but despite being monitored with the same temporal coverage, no planet was detected around \\cyga.  Continued radial-velocity monitoring has yielded no additional planet signatures for either star (D. Fischer, private communication).  Imaging observations, however, do indicate that \\cyga\\ has a faint M dwarf binary companion with a separation of $\\sim 3$'', corresponding to projected separation of $\\sim 70$ AU at the measured distance of the system \\citep[$\\sim 22$ pc;][]{1999PASP..111..321H,2001AJ....121.3254T,2002ApJ...581..654P}. Whether these two objects are gravitationally bound has yet to be determined firmly, but initial proper motion measurements do suggest that they are physically associated \\citep{2002ApJ...581..654P}.  \\citet{1993A&A...274..825F} and subsequently \\citet{1997AJ....113.1871K} found that \\cygab\\ differ in another fundamental way: their Li abundances.  The photospheric Li abundance of \\cygb\\ is a factor $\\geq 4.5$ lower than that of \\cyga.  While both stars are depleted in Li relative to the Solar System's meteoritic value \\citep[$\\log N(\\mathrm{Li})=3.26$;][]{2009ARA&A..47..481A}, the Li abundance of \\cyga\\  ($\\log N(\\mathrm{Li})=1.27$) is slightly higher and that of \\cygb\\ is lower ($\\log N(\\mathrm{Li}) \\leq 0.60$) than that of the Sun ($\\log N_{\\odot}(\\mathrm{Li})=1.05$) \\citep{1997AJ....113.1871K}.  The difference in the Li abundances of \\cygab\\ cannot be explained by standard stellar models, which predict Li depletion is a function of stellar age, mass, and composition; empirical evidence suggests that an extra parameter is needed. \\citet{1997AJ....113.1871K} argue that a slow mixing mechanism, possibly related to rotation, can account for the low absolute Li abundances of both stars, and they discuss a possible connection between Li depletion and planet formation as an explanation for the difference between the two.  More recently, others have also argued that an extra parameter (beyond standard models) is needed to account for the observed Li abundances of solar-type stars \\citep[e.g.,][]{2008A&A...489..677P}.  \\citet{2000AJ....119.2437D} note that the Li--\\teff\\ trend could be quite steep for solar twins, consistent with the 16 Cyg A -- Sun -- 16 Cyg B pattern, so that even if initial angular momentum ($J_{\\mathrm{o}}$) and rotational history do play the role of the extra parameter, $J_{\\mathrm{o}}$ need not be unreasonably different between A and B. \\citet{2000AJ....119.2437D} also found that the Be abundances of \\cygab\\ are the same within the measurement uncertainties, placing an additional constraint on the mechanism responsible for the disparate Li abundances.  In this Letter we present the results of a detailed abundance analysis of 15 elements of the solar twins \\cyga\\ (HR\\,7503, HD\\,186408, HIP\\,96895) and \\cygb\\ (HR\\,7504, HD\\,186427, HIP\\,96901) based on high-resolution echelle spectroscopy.  The abundances allow us to constrain more fully the physical similarities of the two stars, and the implications for Li depletion and planet formation in this system are discussed. ",
        "conclusions": "\\label{s:con} We have presented the results of a detailed abundance analysis of the solar twins \\cygab, the second of which is host to a giant planet.  Aside from a factor of $\\sim 4.5$ difference in Li abundances, the two stars are found to be otherwise chemically identical based on the 15 elements considered.  Slopes in the [m/H]-\\tc\\ relations are also statistically identical and are another indication that \\cygab\\ are chemically homogeneous.  The stark consistency of the compositions of these stars suggest that the physical process(es) responsible for the enhanced Li depletion in B did not alter the abundances of other elements. This argues against any kind of accretion related mechanism and supports differences in internal mixing efficiencies possibly related to different angular momentum evolutions as the most likely explanation for the disparate Li abundances.  Enhanced Li depletion in B can be plausibly tied to the presence of its giant planet, as predicted by rotational stellar evolution models; however, the mixed observational results regarding Li abundances of planet host stars cloud this issue.  More work is clearly required to understand how star-disk interactions and/or planet formation does or does not increase Li depletion in planet host stars.  The chemical homogeneity of \\cygab, combined with the heretofore lack of detected planets around \\cyga, further suggests that the planet formation process did not affect the bulk composition of \\cygb.  Since the discovery that stars with giant planets tend to be more metal-rich than stars without known planets \\citep{1997MNRAS.285..403G,1998A&A...334..221G,2004A&A...415.1153S,2005ApJ...622.1102F}, countless abundance studies of host stars have aimed to identify possible chemical vestiges of the planet formation process.  As described above, Li may be one of these.  As for the overall metallicity of planet hosts, the result for \\cygab\\ adds to the considerable evidence indicating that the planet-metallicity correlation for stars with giant planets is intrinsic in nature and does not arise from processes, such as accretion of solid-body material, associated with the formation and evolution of giant planets.  Furthermore, it appears that the abundances of individual elements heavier than Li \\citep[with the possible exception of Be and B, the abundances of which can also be depleted by internal mixing mechanisms, depending on the depth and efficiency of the mixing;][]{1998ApJ...498L.147D,2005ApJ...621..991B} are also not affected by planet formation, at least in systems like \\cygb.  The physical characteristics of 16 Cygni make it an ideal system to test and constrain planet formation models.  Most tellingly, the conditions necessary for planet formation apparently were present for \\cygb\\ but not \\cyga, despite their physical and chemical similarities.  We have discussed empirical and computational results that can possibly account for the observed characteristics of the system, including the lack of a detected planet around \\cyga, the enhanced Li depletion of \\cygb, and the eccentricity of the planet \\cygb\\ b, and that imply that neither \\cyga\\ nor \\cygb\\ is a terrestrial planet host.  Future efforts that can combine all of these attributes into a single model will represent a significant achievement in understanding the formation and evolution of planetary systems."
    },
    "1107/1107.1531.txt": {
        "abstract": "Integral field spectroscopy of the LV\\,2 proplyd is presented taken with the VLT/FLAMES Argus array at an angular resolution of 0.31$\\times$0.31 arcsec$^2$ and velocity resolutions down to 2 km\\,s$^{-1}$ per pixel. Following subtraction of the local M42 emission, the spectrum of LV\\,2 is isolated from the surrounding nebula. We measured the heliocentric velocities and widths of a number of lines detected in the intrinsic spectrum of the proplyd, as well as in the adjacent Orion nebula falling within a 6.6 $\\times$ 4.2 arcsec$^2$ field of view. It is found that far-UV to optical collisional lines with critical densities, \\crd, ranging from 10$^3$ to 10$^9$ \\cmt\\ suffer collisional de-excitation near the rest velocity of the proplyd correlating tightly with their critical densities. Lines of low \\crd\\ are suppressed the most. The bipolar jet arising from LV\\,2 is spectrally and spatially well-detected in several emission lines. We compute the \\foiii\\ electron temperature profile across LV\\,2 in velocity space and measure steep temperature variations associated with the red-shifted lobe of the jet, possibly being due to a shock discontinuity. From the velocity-resolved analysis the ionized gas near the rest frame of LV\\,2 has \\elt\\ $=$ 9200 $\\pm$ 800\\,K and \\eld\\ $\\sim$ 10$^6$\\,\\cmt, while the red-shifted jet lobe has \\elt\\ $\\approx$ 9000 -- 10$^4$\\,K and \\eld\\ $\\sim$ 10$^6$ -- 10$^7$\\,\\cmt. The jet flow is highly ionized but contains dense semi-neutral clumps emitting neutral oxygen lines. The abundances of \\np, \\opp, \\nepp, \\fepp, \\sulp, and \\sulpp are measured for the strong red-shifted jet lobe. Iron in the core of LV\\,2 is depleted by 2.54 dex with respect to solar as a result of sedimentation on dust, whereas the efficient destruction of dust grains in the fast microjet raises its Fe abundance to at least 30 per cent solar. Sulphur does not show evidence of significant depletion on dust, but its abundance both in the core and the jet is only about half solar.   %\\hep, \\cpp\\ and \\opp\\ are derived using recombination lines, and of \\np, \\opp, %\\nepp, \\fepp using collisionally excited lines and compared the kinematics of %several ionic species. ",
        "introduction": "Protoplanetary disks (proplyds) provided the first evidence of gaseous dusty disks around young stars in the early 1990s. The archetypal proplyds found in the Orion nebula (M42; O'Dell et al. 1993) are dense semi-ionized objects photoevaporated to various degrees depending on their distance from the ionizing Trapezium cluster. They are landmark objects in the study of how circumstellar disks and eventually planetary systems form (e.g. O'Dell 2001). They are a unique environment for the study of disk formation and evolution in areas dominated by massive OB-type stars. Massive stellar clusters and their associated \\hii\\ regions, such as in the Orion OB1 association, are thought to represent the closest analogs to the birth environment of our Solar System (Adams 2010). The proplyds of Orion should therefore be ideal laboratories for the study of our planetary system's origins. At optical wavelengths proplyds present a photoionized skin facing the ionizing cluster, giving way to a dusty envelope often shaped into comet-shaped outflows. The elemental content and chemistry of proplyds are \\emph{very poorly} known, however, (i) the composition of planet-forming circumstellar envelopes is of great topical interest given the established positive correlation between host star metallicity and the incidence of giant planetary companions (e.g. Gonzalez 1997; Neves et al. 2009); (ii) higher metallicity in the protoplanetary disk favours the formation of giant planets in the `core accretion' scenario (e.g. Boss 2010). While the previously discrepant abundances for the interstellar medium and main-sequence stars in Orion are beginning to converge (Sim\\'on-Diaz \\& Stasinska 2011), proplyds are the last major component of Orion for which a chemical abundance scale is lacking.  We have been taking steps to rectify this situation with a programme targeting a sample of bright proplyds in M42. In Tsamis et al. (2011; Paper~I hereafter) a chemical abundance study of the protoplanetary disk LV\\,2 (Laques \\& Vidal 1979) and its Orion nebula host vicinity was presented, based on the analysis of VLT optical integral field spectroscopy and \\hst\\ FOS single aperture ultraviolet to far-red spectroscopy. From an emission-line analysis the abundances of several elements were measured for the proplyd and the local M42 nebular field. LV\\,2 was found to be slightly overabundant in carbon, oxygen and neon compared to the Orion nebula gas-phase composition and to the Sun. The carbon, oxygen and neon abundances in LV\\,2 were measured to be $\\approx$0.2--0.3 dex higher than those in B-type stars of the Ori OB1 association studied by Sim\\'{o}n-Diaz (2010) and Sim\\'on-Diaz \\& Stasinska (2011).  That result constitutes a direct measure of the metallicity of gas photo-evaporated from circumstellar disk material where planet formation may be underway.  In this paper, we present higher velocity dispersion integral field spectra of LV\\,2 which enable a physical analysis of its fast microjet to be undertaken. The paper is organized as follows. The dataset specifications and reduction method are outlined in Section 2. The analysis and results are presented in Section 3, followed by our conclusions in Section~4. ",
        "conclusions": "We have presented spatially resolved, high dispersion spectra of LV\\,2 and its Orion nebula vicinity based on VLT observations. The deep spectra enabled the detection and resolution of emission lines arising from the proplyd's irradiated bipolar microjet. Maps of LV\\,2 in the light of \\hi\\ and \\ffeiii\\ lines have been presented, where the jet lobes appear spatially resolved from the core of the proplyd.  An abundance analysis was undertaken for the brighter red-shifted jet lobe complementing the abundance analysis of the core of LV\\,2 presented in Paper~I. A noteworthy result is the significantly enhanced abundance of iron in the jet versus the core of LV\\,2, pointing to very efficient dust destruction mechanisms in fast irradiated microjets. The overall depletion of iron in the proplyd's rest frame is higher than in Galactic \\hii\\ regions and is comparable to what is found along cold ISM sight lines. The Ne/S abundance ratio in LV\\,2 as probed by optical lines is $\\geq$2 times solar, in agreement with the value applicable to the generic M42 gas from IR line measurements by {\\it Spitzer}.  The extreme-UV irradiated jet shows a range of ionization conditions allowing the survival of neutral oxygen pockets mixed in with highly ionized gas. A velocity-resolved analysis was performed probing the electron temperature and density across LV\\,2's rest frame and red-shifted jet lobe. The \\foiii\\ and \\fciii\\ emitting regions in LV\\,2 show a range of densities with a peak of 10$^6$ \\cmt\\ near the line systemic velocity and an upper limit of $\\sim$10$^7$ \\cmt\\ from \\ffeiii\\ emission within the red-shifted jet lobe. The electron temperature in the highly ionized zone in the main body of the proplyd shows a small 10 per cent rms variation. Larger temperature variations associated with the onset of the jet emission have been attributed to a shock discontinuity.  The far-UV to optical collisional lines emitted from LV\\,2 are correlatively quenched according to the critical densities of their upper levels, with lines of low critical density being affected the most. This reinforces the conclusion of Paper~I that in \\hii\\ regions containing dense zones, such as proplyds, filaments, globules etc., line ratios such as \\foiii\\ $\\lambda$4363/($\\lambda$4959$+$$\\lambda$5007) and \\fnii\\ $\\lambda$5755/($\\lambda$6548$+$$\\lambda$6584) will be heavily weighted towards the dense component of the gas. These line ratios are the main temperature diagnostics for Galactic and extragalactic \\hii\\ regions, whose validity rests on an independent way of determining the mean density of the plasma. This is usually done via the \\foii\\ or \\fsii\\ doublet ratios which are demonstrably biased towards low density regions (Fig.\\,~5 left). The \\fnii\\ ratio will be especially biased as the critical densities of its constituent lines are at a ratio of approximately 200 versus only 40 for the \\foiii\\ lines, resulting in more uneven fractions of \\fnii\\ $\\lambda$5755 versus \\fnii\\ $\\lambda$6584 to be emitted from dense and diffuse gas respectively, than is the case for \\foiii\\ $\\lambda$4363 and \\foiii\\ $\\lambda$5007 -- this is demonstrated in Fig.\\,~5.  In this situation, and especially in cases of distant nebulae whose gas density distribution appears deceptively homogeneous on large scales, or in the presence of condensations just below the angular resolution limit in nearby nebulae, the end result would be an overestimation of the electron temperature and an inevitable underestimation of the gas metallicity (e.g. Rubin 1989). Given the behaviour of the \\fnii\\ and \\foiii\\ lines seen in Fig.\\,~5 the former temperature would be more adversely affected that the latter in \\hii\\ regions with a clumpy small-scale density distribution. It may be no coincidence therefore that \\elt(\\fnii) values in \\hii\\ regions are typically higher than \\elt(\\foiii) values by $\\simeq$1000\\,K on average (e.g. Rodr\\'iguez \\& Delgado-Inglada 2011). This is usually justified in the literature as being due to the fact that the low ionization zones in a nebula tend to have a higher temperature than the \\oiii\\ zone. Based, however, on the arguments presented here this behaviour could be at least partly spurious if the underlying density distribution of a nebula is sufficiently clumpy at sub-arcsecond scales.  Regarding the abundance anomaly pertaining to RL versus CL diagnostics (e.g. Tsamis et al. 2003; Tsamis \\& P\\'equignot 2005; Garc\\'ia-Rojas \\& Esteban 2007; Ercolano 2009; Esteban et al. 2009; Mesa-Delgado et al. 2010; Sim{\\'o}n-D{\\'{\\i}}az \\& Stasi{\\'n}ska 2011; Paper~I), it should be noted that the fraction of \\oii\\ $\\lambda$4649, \\cii\\ $\\lambda$4267 and \\hi\\ RLs that are emitted from LV\\,2 is the same for all these lines ($\\approx$70 per cent), while the rest is emitted from diffuse M42 gas along the line of sight. As a result the proplyd does not show elevated \\opp/\\hp\\ and \\cpp/\\hp\\ abundance ratios, measured using RLs, compared to the Orion background (cf. Paper~I). In our view this analysis goes some way towards explaining the abundance anomaly in \\hii\\ regions as a problem affected by the use and interpretation of the \\foiii\\ and \\fnii\\ plasma thermometers when applied to average spectra of large nebular volumes."
    },
    "1107/1107.4289_arXiv.txt": {
        "abstract": "We describe a new algorithm for finding substructures within dark matter haloes from N-body simulations. The algorithm relies upon the fact that dynamically distinct substructures in a halo will have a {\\em local} velocity distribution that differs significantly from the mean, {\\em i.e.} smooth background halo. We characterize the large-scale mean field using a coarsely grained cell-based approach, while a kernel smoothing process is used to determined the local velocity distribution. Comparing the ratio of these two estimates allows us to identify particles which are strongly cluster in velocity space relative to the background and thus resident in substructure. From this population of outliers, groups are identified using a Friends-of-Friends-like approach. False positives are rejected using Poisson noise arguments. This approach does not require a search of the full phase-space structure of a halo, a non-trivial task, and is thus computationally advantageous. We apply our algorithm to several test cases and show that it identifies not only subhaloes, bound overdensities in phase-space, but can recover tidal streams with a high purity. Our method can even find streams which do not appear significantly overdense in either physical or phase-space. ",
        "introduction": "\\label{sec:intro} The dark matter in the haloes that surround galaxies such as our own is often modeled as a smooth distribution in both position and velocity space. In particular, the spatial distribution is generally assumed to fall off monotonically from the halo center, and the local velocity distribution is assumed to be Maxwellian. These assumptions provide a useful starting point for making model predictions for dark matter detection experiments ({\\em e.g.}~\\citealp{lewin1996}). However, the haloes that form in cosmological simulations are anything but smooth, even relaxed, virialized haloes that have not undergone any major mergers in recent history. In particular, they contain subhaloes, which range in mass from a few percent of the host halo's mass down to the resolution limit of the simulation ({\\em e.g.}~ \\citealp{moore1999,gao2004,diemand2006,diemand2008,springel2008}).  \\par Subhaloes are but one example of substructure found in simulated haloes. As subhaloes traverse the host they will experience mass loss and leave behind debris in the form of shells and streams. In simulations these coherent streams of particles can exist for a few orbital periods before eventually dissolving into the smooth halo background. The resulting superposition of streams and subhaloes produces a complex velocity structure (e.g.~\\citealp{fairbairn2009,kuhlen2010}), and even the ``smooth'' background of a halo is composed of many overlapping fundamental dark matter streams \\citep{stiff2003,vogelsberger2010}. Lastly, and perhaps most practically, the number, velocity profile and mass of these streams will have important ramifications for direct and indirect dark matter searches ({\\em e.g.}~\\citealp{stiff2001,vogelsberger2008,vogelsberger2009,kuhlen2010}).  \\par The interest in streams is not limited solely to dark matter. Galactic stellar haloes are thought to form, in part, from the accretion and subsequent disruption of satellite galaxies \\citep{johnston2008}. A prime example is the Sagittarius stream associated with the Sagittarius dwarf \\citep{ibata1994,ibata2001}. The PANDAS survey has found evidence of stellar streams in M31, which may come from disrupted satellites \\citep{pandas2009}.  These stellar substructures offer an observable window into galaxy formation history and searches for such structures are just beginning in earnest. Data from future instruments such as GAIA \\citep{gaia2008} will allow us to test the $\\Lambda$CDM model by comparing the number of observed stellar streams to the number predicted from cosmological simulations.  \\par The problem with finding substructures is non-trivial and a number of techniques have been developed to find subhaloes. These include {\\sc skid} \\citep{skid}, {\\sc subfind} \\citep{subfind}, \\6DFOF\\ \\citep{diemand2006}, \\enlink\\ \\citep{enlink}, and \\hsf\\ \\citep{hsf}. Algorithms such as \\citetalias{subfind} search for physical overdensities or clustering in configuration space and are well suited to finding subhaloes. Simulations show that the contrast between substructures and the background is more prominent in phase-space and action space \\citep{helmi2000,gomez2010}. Algorithms such as \\citetalias{enlink} and \\citetalias{hsf} make use of this increased contrast by using phase-space density to identify substructures. In general, most of these algorithms search for local peaks in either physical or phase-space density and estimate the outer boundaries enclosing these peaks. A number of studies have searched for completely disrupted subhaloes by identifying particles that were once bound to a subhalo at some early time and that are no longer bound to this subhalo \\citep{helmi2003,vogelsberger2009b,xu2009}. This process is laborious and computationally intensive since past snapshots must be stored and tracked from output to output.  \\par Here we present a new algorithm, the STructure Finder (\\STF) designed to find velocity substructures within dark matter haloes. Substructures occupy specific regions in orbital space and have smaller velocity dispersion than the virialized background and consequently should appear more strongly clustered in velocity space relative to the background. In short, we identify candidate streams (and subhaloes) by finding groups of particles that \"stand out\" against the velocity distribution of the background. We model the local velocity distribution as a multivariate Gaussian and therefore, our method is most sensitive to substructures on the tails of the Gaussian, that is, substructures that are cold but that have a large bulk velocity. Once we have identified these outlying particles we then use a Friends-of-Friends-like algorithm to link them.  \\par Our paper is organized as follows: In section \\Secref{sec:algorithm}, we outline our algorithm, determine the algorithm's optimal parameters, and discuss what types of substructure can be identified. We test our algorithm with toy models and discuss the results in \\Secref{sec:toymodels}. In section \\Secref{sec:halomany}, we analyze a more realistic model of a dark matter halo. The paper concludes in \\Secref{sec:discussion} with a summary and discussion. ",
        "conclusions": "\\label{sec:discussion} In this paper, we present an algorithm designed to identify both tidal streams and subhaloes.  We begin by dividing the halo into cells. In each cell, we model the velocity distribution function of the background as a multivariate Gaussian (the ``Maxwellian sea'').  While there is no exact method for identifying and modelling the background population, we have demonstrated that decomposing the halo into sub-volumes (cells) containing a relatively large number of particles is sufficient to average out any substructures present. We then calculate an effective velocity-space density for each particle in the cell. The ratio of the {\\em local} velocity density to the density of the background at the particle's velocity provides a probabilistic measure of whether the particle belongs to the background or to a stream or subhalo. This key step of identifying outliers is then combined with a FOF-like search to identify substructures.  \\par We have clearly shown using several test cases that our method is capable of identifying tidal streams and any subhaloes present in a halo. The purity of the identified substructures is very high regardless of the type of substructure, whether subhalo or diffuse tidal stream. Our method is even capable of identifying streams with a high purity in cases where the particles of the stream do not appear significantly overdense in either physical or phase-space for a given phase-space estimator, and account for a small fraction of the mass in the spatial volume occupied by the stream. Furthermore, we are able to effectively identify substructures in the cores of haloes, even haloes with low resolution.  \\par We also note that our key step of identifying velocity outliers can also be applied to any quantity for which one can characterize the mean distribution. The increased contrast between particles belonging to substructures and those belonging to the  background makes searching the resulting subset simpler. Furthermore, one could combine this key step with a variety of searches tailored to identify specific substructures.  \\par Our algorithm is not the only one capable of finding tidal streams. Both \\hsf\\ \\citep{hsf} and \\enlink\\ \\citep{enlink} should be able to find tidal streams, however, both these algorithms rely on accurately estimating the phase-space density for each particle. In general, determining the velocity density is less computationally intensive than accurately computing the phase-space density. Searching the phase-space of a halo requires methods that calculate the local six dimensional metric for each particle to accurately determine distances in phase-space. We have also shown that, in certain cases, identifying substructures solely based on the phase-space density may miss diffuse tidal streams.  \\par The presence of tidal streams in the central regions of haloes has potentially important ramifications for indirect and direct dark matter searches. For instance, most studies of direct dark matter detection assume the velocity distribution of dark matter in the solar neighbourhood of a galactic mass halo is Maxwellian in order to determine the detection rate and recoil spectrum. Simulations show that the velocity distribution is more complex, possibly due to the presence of substructure. A few studies have attempted to go beyond this simple assumption but are limited by the resolution of the studies used to determine the velocity distribution \\citep{fairbairn2009, kuhlen2010}. Since our algorithm allow us to accurately characterize the stream distribution function with a single snapshot from a cosmological simulation, it provides a easily accessible avenue for going beyond the resolution of current simulations. By searching numerous haloes from numerous cosmological simulations and stacking results, we will be able to determine the distribution function of these streams and extrapolate it to improve the accuracy of detection rates and the observed energy spectrum."
    },
    "1107/1107.1715_arXiv.txt": {
        "abstract": "{ We present updated global fits of the constrained Minimal Supersymmetric Standard Model (cMSSM), including the most recent constraints from the ATLAS and CMS detectors at the LHC, as well as the most recent results of the XENON100 experiment. Our robust analysis takes into account both astrophysical and hadronic uncertainties that enter in the calculation of the rate of WIMP-induced recoils in direct detection experiment. We study the consequences for neutralino Dark Matter, and show that current direct detection data already allow to robustly rule out the so-called Focus Point region, therefore demonstrating the importance of particle astrophysics experiments in constraining extensions of the Standard Model of Particle Physics. We also observe an increased compatibility between results obtained from a Bayesian and a Frequentist statistical perspective. We find that upcoming ton-scale direct detection experiments will probe essentially the entire currently favoured region (at the 99\\% level), almost independently of the statistical approach used. Prospects for indirect detection of the cMSSM are further reduced.} ",
        "introduction": "\\label{secintro} The constrained Minimal Supersymmetric Standard Model (cMSSM) \\cite{Chamseddine:1982jx,Kane:1993td} is the most widely discussed extension of the Standard Model (SM) of particle physics. Despite its relative simplicity, this model has the advantage of capturing some key phenomenological feature of Supersymmetry (SUSY), while making definite predictions for the properties of the the lightest neutralino $\\tilde{\\chi}_1^0$ (which we denote $\\chi$ here, for simplicity), a linear superposition of the superpartners of the neutral  Gauge bosons and the neutral Higgses, which is by far the most popular Dark Matter (DM) candidate \\cite{Bertone:2010zz,Bergstrom:2000pn,Munoz:2003gx,Bertone:2004pz}.  Recently, the ATLAS and CMS experimental collaborations have published the first search for supersymmetry (SUSY) \\cite{Chatrchyan:2011wc,Khachatryan:2011tk,daCosta:2011qk,Aad:2011ks,Aad:2011xm}, based on  $\\sqrt{s} = 7$ TeV proton-proton collisions at the LHC recorded during 2010 (see the experimental papers above for details on the event topologies, selection cuts, etc.). Since no excess above the SM predictions has been observed, these searches have set new interesting constraints on physics beyond the SM, and several groups of authors have already studied the impact of these results on the cMSSM and more in general on SUSY (see e.g. Refs. \\cite{Buchmueller:2011aa,Allanach:2011ut,Akula:2011dd,Strumia:2011dv,Allanach:2011wi,Buchmueller:2011ki,Farina:2011bh,Strumia:2011dy}).  Here, we perform new global fits of the cMSSM, following the methodology outlined in Refs. \\cite{deAustri:2006pe,Trotta:2008bp,Feroz:2011bj}, where LEP constraints, precision tests of the SM and the most recent cosmological constraint on the relic abundance of the DM were implemented using nested sampling as a scanning algorithm, in order to derive both the most probable regions of the cMSSM (Bayesian) and the best fit regions (frequentist). We update the analysis to include new $B$-physics observables, as well as both the aforementioned LHC constraints and the upper limit on the WIMP-nucleon scattering cross-section arising from DM direct detection experiments (see e.g. Ref.~\\cite{Cerdeno:2010jj} and references therein). We implement the findings of the XENON100 collaboration \\cite{xenon:2011hi}, which recently reported the results of a two-phase time projection chamber with 48 kg of liquid xenon as ultra-low background fiducial target. In 100.9 live days of data, acquired between January and June 2010, the collaboration recorded three candidate events, with an expected background of 1.8 $\\pm$ 0.6 events. The rate of DM-induced recoil events in a direct detection experiment being directly proportional to the spin-independent scattering cross section of neutralinos off protons $\\sigmaSI$, the upper limit on the recoil rate translates into an upper limit on $\\sigmaSI$.  The predictions for $\\sigmaSI$ within the cMSSM (given constraints other than direct detection data) span approximately the range between $10^{-7}$ pb to $10^{-10}$ pb, and since XENON100 excludes  $\\sigmaSI$ larger than $7.0 \\times 10^{-9}$ pb for a DM particle mass of 50 GeV at the 90\\% confidence level, one sees immediately that this has the potential to rule out significant portions of the cMSSM, under the hypothesis of a standard expansion rate of the Universe, and that the neutralinos make up most of the DM in the Universe.  Particular care must be paid, however, when translating the information provided by a direct detection experiment to constraints on the cMSSM parameter space, since this requires specific modeling of the velocity distribution of DM particles and of the DM density in the solar neighborhood. As a consequence, the upper limit from XENON100 can be considered robust only if the uncertainties on these quantities are taken into account. Moreover, the value of $\\sigmaSI$ depends on nuclear physics quantities, namely the hadronic matrix elements, which parametrize the quark composition of the proton, and this affects the way direct detection data are applied to constrain cMSSM parameters. Therefore in the present analysis we incorporate the local density, the quantities defining the DM velocity distribution and those modeling the proton composition as nuisance parameters, which are then either marginalized (i.e., integrated out) or maximised over to obtain inferences on the cMSSM parameters that fully account for those astrophysical and hadronic uncertainties.  The paper is organized as follows: in Section~\\ref{sec:theory} we present our theoretical framework, in Section~\\ref{secDD} we present the formalism for the computation of recoil events, our modelization of the DM velocity distribution and our implementation of the XENON100 data. Section \\ref{secresults} is devoted to the presentation of the results, while prospects for the discovery of SUSY in the cMSSM are discussed in Section \\ref{sec:prospects}. We conclude in Section \\ref{secconclusion}. ",
        "conclusions": "\\label{secconclusion}   We have presented in this paper new global fits of the cMSSM, including the most recent constraints from the LHC and the XENON100 experiment. Besides the uncertainties on Standard Model quantities, our analysis takes into account both astrophysical and hadronic uncertainties that enter in the calculation of the rate of WIMP-induced recoils in direct detection experiment. We have shown that the FP branch of the cMSSM is robustly ruled out, even when all uncertainties are taken into account. These results highlight the complementarity of direct detection experiments to collider searches. Indeed, direct detection data would still be required to rule out the FP even if the LHC exclusion limits are extended to a total integrated luminosity of 7 fb$^{-1}$ at 7 TeV center-of-mass energy~\\cite{Bechtle:2011dm}, which might be achieved by the end of 2012.  Although our results are specific to the cMSSM, we note that the conditions for the occurrence of the FP can actually be extended to a general class of unconstrained Supersymmetric scenarios. In models where no universality assumptions on the soft mass parameters are made, the FP region can be shifted to either larger or lower scalar mass parameters, but the general feature of the smallness of the $\\mu$ parameter is preserved, as are its implications for dark matter. In Ref.~\\cite{Baer:2005ky}, for example, a low-energy parametrization of the FP was studied that could be used to describe this region irrespectively of the high-energy assumptions. Although the analysis of these general constructions is beyond the scope of the present work, it seems reasonable to extend our conclusions to the general case, given the similarity of the predicted direct detection rates.  While this work was under completion, Refs.~\\cite{Farina:2011bh,Buchmueller:2011ki} performed a similar analysis of the impact of LHC and XENON100 data on the parameter space of several theoretical models, including the cMSSM. While our results about the viability of the FP regions are qualitatively similar, our analysis differs from that in the above references in that we studied the impact of both hadronic and astrophysical uncertainties (both neglected in Ref.~\\cite{Farina:2011bh}, while Ref.~\\cite{Buchmueller:2011ki} includes only hadronic uncertainties). Furthermore, we adopt a more sophisticated statistical framework than Ref.~\\cite{Farina:2011bh} (where the parameter space was explored using random scans, rather than the more sophisticated statistical scans employed here) and perform a detaild, quantitative comparison between the Bayesian and profile likelihood results (absent in Ref.~\\cite{Buchmueller:2011ki})."
    },
    "1107/1107.2185_arXiv.txt": {
        "abstract": "We propose a new method to explore the candidate super-Eddington active galactic nuclei (AGNs). We examine the properties of infrared (IR) emission from the inner edge of the dusty torus in AGNs, which are powered by super- or sub-Eddington accretion flows around black holes, by considering the dependence of the polar angle on the radiation flux of accretion flows \\citep{Wa05}. We find that for super-Eddington AGNs, of which the mass accretion rate is more than $10^{2}$ times larger than the Eddington rate, the ratio of the AGN IR luminosity and the disc bolometric luminosity is less than $10^{-2}$, unless the half opening angle of the torus ($\\theta_{\\rm torus}$) is small ($\\theta_{\\rm torus}<65^\\circ$). This is due to the self-occultation effect, whereby the self-absorption at the outer region of the super-Eddington flow dilutes the illumination of the torus. Such a small luminosity ratio is not observed in sub-Eddington AGNs, whose mass accretion rate is comparable to or no more than 10 times larger than the Eddington mass accretion rate, except for extremely thin tori ($\\theta_{\\rm torus}>85^\\circ$). We also consider the properties of the near-IR (NIR) emission radiated from hot dust $>$ 1000 K. We find that super-Eddington AGNs have a ratio of the NIR luminosity to the bolometric luminosity, $L_{\\rm NIR,AGN}/L_{\\rm bol,disc}$, at least one order of magnitude smaller than for sub-Eddington AGNs for a wide range of half opening angle ($\\theta_{\\rm torus} > 65^{\\circ}$), for various types of dusty torus model. Thus, a relatively low $L_{\\rm NIR,AGN} /L_{\\rm bol,disc}$ is a property that allows identification of candidate super-Eddington AGNs. Lastly, we discuss the possibility that NIR-faint quasars at redshift $z\\sim 6$  discovered by a recent deep SDSS survey may be young quasars whose black holes grow via super-Eddington accretion. ",
        "introduction": "It is now widely accepted that most normal galaxies have central supermassive black holes (SMBHs) with $M_{\\rm BH} =10^{6\\rm -9}M_{\\odot}$ (e.g.,\\citealt{KR95,Mi95,Gi09}). High resolution observations of nearby galaxies have revealed a close correlation between the mass of SMBHs and the bulge mass, or the velocity dispersion of the bulge (e.g., \\citealt{KR95,FM00,MH03,HR04}). This implies that the formation of SMBHs is strongly coupled with the formation of galaxies. However, the formation and evolution of SMBHs is still an open question and a hot topic in astrophysics.  The discovery of luminous quasars at redshift $z >6$  shows that SMBHs of the order of billions of solar masses exist at the end of the re-ionization epoch (e.g., \\citealt{Fa01,Fa06,Go06,Wi07,Wi10}). Considering Eddington-limited accretion, $z > 6$ quasars may not form readily. Unless the seed black hole mass is large, gas accretion with low radiative efficiency and/or BH mergers significantly contribute to the growth of SMBHs (see \\citealt{Sh05}). If super-Eddington accretion is possible, the growth timescale of BHs can be much shorter than the Eddington timescale (e.g., \\citealt{Ka04,Oh05}). Thus, super-Eddington accretion might explain the formation of SMBHs. Recently, \\citet{KW09} predicted that super-Eddington accretion is required for the formation of $z>6$ quasars with $M_{\\rm BH}\\approx 10^{9}M_{\\odot}$, because the final BH mass is greatly suppressed by star formation in a circumnuclear disc and AGN outflow based on the coevolution model of SMBH growth and a circumnuclear disc (\\citealt{KW08}; see also \\citealt{VR05}).  The occurrence of super-Eddington accretion flow has not yet been accurately verified, although this issue has been investigated since the 1970s (e.g., \\citealt{SS73,Be78,Ab88,HC91,WF99,Wa00}. By using two-dimensional radiation hydrodynamic simulations, \\citet{Oh05} demonstrated for the fist time that quasi-steady super-Eddington disc accretion is possible, because both the radiation field and the mass accretion flow around a BH are non-spherical and the large number of photons generated in the disc is swallowed by the BH without being radiated away due to photon-trapping (see \\citealt{OM07}). This conclusion has recently been confirmed by two-dimensional radiation magnetohydrodynamic simulations \\citep{Oh09}. However, it is still unclear whether super-Eddington accretion can significantly contribute to the formation of SMBHs, because the radiation pressure and/or disc wind associated with super-Eddington accretion might play an important role in the gas accretion process in host galaxies (e.g., \\citealt{SR98,Fa99,Ki03,Oh07}).  In the local Universe, we have discovered a class of AGNs with a higher Eddington ratio among the AGN population, i.e., Narrow Line Seyfert 1 galaxies (NLS1s). They have characteristic properties such as narrow Balmer lines (e.g., \\citealt{OP85,Po00}), strong soft X-ray excess (e.g., \\citealt{Po95,Bo96,Bo03}) and rapid variability (e.g., \\citealt{Ot96,Le99,Bo02,Ga04}). These properties indicate that they have a small black hole mass and high Eddington ratio (e.g., \\citealt{BB98,Ha00,Mi00}). The observed bolometric luminosity of NLS1s saturates at a few times the Eddington luminosity (e.g., \\citealt{CK04,Co06,Mu08,La09}), which is consistent with predictions of the super-Eddington accretion disc (e.g., \\citealt{Wa00,Oh05,Oh09}). In addition, the star formation rate of NLS1 hosts is higher than that of BLS1s (\\citealt{Sa10}). Thus, NLS1s may be the early phase of rapid BH growth via super-Eddington accretion (e.g., \\citealt{Ma00,Ka04}). In the high-$z$ universe, we may speculate that super-Eddington AGNs are a more dominant population of AGNs because the lifetime of SMBH growth is closer to the cosmic age if the formation of SMBHs is mainly a gas accretion process. To confirm this, it is important to determine whether the fraction of super-Eddington AGNs increases with redshift. Some observations to date have indicated that the average Eddington ratio increases  slightly with redshift (e.g., \\citealt{Ko06,Sh08,Wi10}). However, we have not observed super-Eddington AGNs among high-$z$ quasars (e.g., \\citealt{MD04,Sh08}) and AGNs in high-$z$ massive galaxies (e.g., \\citealt{Ya09}). Thus, it is essential to find candidate super-Eddington AGNs before starting detailed observations.  According to unification models of AGNs, the accretion disc is surrounded by a dusty structure (e.g., \\citealt{An93,El08}). A significant fraction of the emitted ultraviolet and optical radiation of the accretion disc is absorbed by the dust and is re-emitted at IR wavelengths. In particular, the hot dust is directly heated by the central engine and produces near-IR (NIR) emission (e.g., \\citealt{Ri78,Ha03}). Since it has been reported that the polar angle dependence of the radiation flux for super-Eddington accretion flow deviates from that for sub-Eddington flow (e.g., \\citealt{Fu00,Oh05,Wa05} hereafter W05), the reprocessed IR emission may be useful for exploring candidates of super-Eddington accretion flow in AGNs. In this paper, we investigate the properties of IR (as well as NIR) emission from the inner edge of the dusty torus, employing radiative flux from the disc accretion flows by W05, in which the radiation fluxes of sub- and super-Eddington accretion discs are given as functions of the polar angle and the mass accretion rate. ",
        "conclusions": "\\subsection{IR emission from NLR clouds} So far, we have not taken into account the IR emission from narrow-line region (NLR) clouds. However, it is known that dust grains are present in the NLR of Seyferts (e.g., \\citealt{DD88,NL93}). Thus, the dusty clouds in the NLR might be a source of IR emission and hence we must consider the contribution of IR emission from NLR clouds.  NLR clouds are assumed to be distributed in a geometrically thin spherical shell with radius $r_{\\rm NLR}$ in the range $0\\leq \\theta \\leq \\theta_{\\rm torus}$. Assuming the covering factor of the NLR clouds, $f_{\\rm NLR}$, and the optical depth of the NLR clouds, $\\tau_{\\rm NLR}$, the IR luminosity from NLR clouds can be given by \\begin{eqnarray} L_{\\rm IR, NLR}&=&4\\pi r^{2}_{\\rm NLR}f_{\\rm NLR}(1-e^{-\\tau_{\\rm NLR}}) \\int_{0}^{\\theta_{\\rm torus}}F(\\theta)\\sin{\\theta}d\\theta.\\nonumber \\\\ {}&=&f_{\\rm NLR}(1-e^{-\\tau_{\\rm NLR}}) \\left(1-\\frac{L_{\\rm IR,AGN}}{L_{\\rm bol,disc}}\\right), \\end{eqnarray} where $L_{\\rm IR,AGN}/L_{\\rm bol,disc}$ is a function of $\\theta_{\\rm torus}$ (see eqs. (5) and (6)). In Fig. 4, we plot the ratio of the IR luminosity to the bolometric luminosity against $\\theta_{\\rm torus}$ for the NLR clouds, assuming $\\tau_{\\rm NLR}=1$ and $f_{\\rm NLR}=0.1$. Note that the dust extinction of NLR emission is small (i.e., $\\tau_{\\rm NLR} \\leq 1$) for type 1 AGNs \\citep{R00} and the typical covering factor is about $10\\%$ (e.g., \\citealt{NL93,BL05}).  Figure 4 shows that the IR emission from the NLR clouds dominates above $\\theta_{\\rm torus}=75^{\\circ}$ for $\\dot{m}_{\\rm BH}=1,10$, while for $\\dot{m}_{\\rm BH}=10^{2}$ and $\\dot{m}_{\\rm BH}=10^{3}$ the IR emission from NLR clouds is important above $\\theta_{\\rm torus} =60^{\\circ}$ and $\\theta_{\\rm torus}=50^{\\circ}$, respectively. This suggests that for super-Eddington AGNs with $\\dot{m}_{\\rm BH}\\geq 10^{2}$, the IR emission from NLR clouds can dominate the IR emission from the dusty torus. This is because the radiative flux of super-Eddington AGNs is highly collimated compared to sub-Eddington AGNs (see eqs. (2) and (3)). However, we should mention that the NLR size, $r_{\\rm NLR}$, is about $10^{2}$ times as large as the sublimation radius (e.g., \\citealt{Be02,Ne04,BL05}) and the dust temperature of NLR clouds is lower than the hottest dust temperature (i.e., 1500 K). Thus, the IR emission from NLR clouds would contribute mainly MIR (10--30 $\\mu$m) emission (see \\citealt{Sc08,Mor09}). In order to avoid contamination of the IR emission from NLR clouds, the ratio $L_{\\rm NIR,AGN}/L_{\\rm bol,disc}$ can be used to probe an inflated disk due to super-Eddington accretion, as mentioned in $\\S 3.2$.  \\begin{figure} \\begin{center} \\includegraphics[height=6.5cm,clip]{f4.eps} \\end{center} \\caption{ $L_{\\rm IR,AGN}/L_{\\rm bol,disc}$ against $\\theta_{\\rm torus}$ for various $\\dot{m}_{\\rm BH}$. The solid lines show the ratio of the IR luminosity and bolometric luminosity for the NLR clouds, assuming $\\tau_{\\rm NLR}=1$ and $f_{\\rm NLR}=0.1$. The dashed lines show $L_{\\rm IR}/L_{\\rm bol,disc}$ for the dusty torus (see Fig. 2). } \\end{figure}   \\subsection{Contribution of X-ray emission from the corona} We here discuss the contribution of X-ray emission from the corona on the IR emission. For the model of X-ray emission from the corona, we adopt a simple lamp-post geometry model, where the X-ray source is located above the SMBH at some height $h_{\\rm x}$ (e.g., \\citealt{GF91,NK01}). First, we examine the condition where the X-ray emission can reach the torus inner edge for super-Eddington accretion when $\\theta_{\\rm torus}\\geq 45^{\\circ}$. If the light ray, which passes though the maximum thickness of the slim disc $h_{\\rm slim}$ from the X-ray source located at $h_{\\rm x}$, hits the torus inner edge, $r_{\\rm sub}$, a large amount of X-ray emission is absorbed by the dust torus. Assuming $r_{\\rm sub}\\gg r_{\\rm slim}$, this condition can be simply expressed by the following: \\begin{equation} \\frac{h_{\\rm x}}{h_{\\rm slim}} > 1-\\tan(90^{\\circ}-\\theta_{\\rm torus}), \\end{equation} where $r_{\\rm slim}$ is the radius of the slim disc and $h_{\\rm slim}=r_{\\rm slim}$ because the half opening angle of the slim disc is $45^{\\circ}$ (see $\\S 2$). For example, the X-ray emission could produce the IR emission, if $h_{\\rm x} > 0.4h_{\\rm slim}$ and $h_{\\rm x} > 0.8h_{\\rm slim}$ for $\\theta_{\\rm torus}=60^{\\circ}$ and $80^{\\circ}$, respectively. According to the theory of slim discs, $r_{\\rm slim}$ is about $10r_{\\rm g}$ for $\\dot{m}_{\\rm BH}=10^{2}$. Note that for $\\dot{m}_{\\rm BH}=10^{3}$, $r_{\\rm slim}$ is much larger than $10r_{\\rm g}$ (e.g., \\citealt{Wa00}). This may imply that the contribution of X-rays from the corona is not negligible if $h_{\\rm x}$ is larger than $4$--$8r_{\\rm g}$. The lower limit of $h_{\\rm x}$ is comparable to $h_{\\rm x}=6r_{\\rm g}$, which is frequently used to discuss the iron-line reverberation in AGNs (e.g., \\citealt{RB97}).  We next investigate the contribution of X-rays from the corona on the IR emission when $h_{\\rm x} > (4$--$8) r_{\\rm g}$. The ratio of the radiation flux from the accretion disc, $F_{\\rm disc}$, and the corona, $F_{\\rm cor}$, at $\\theta_{\\rm torus}$ is expressed as \\begin{equation} \\frac{F_{\\rm disc}}{F_{\\rm cor}}=\\left(\\frac{L_{\\rm bol,disc}} {L_{\\rm x}}\\right)\\left(\\frac{r_{\\rm x}}{r_{\\rm sub}}\\right)^{2} \\left \\{ \\begin{array}{l} 2\\cos{\\theta_{\\rm torus}}\\,\\,\\, (\\dot{m}_{\\rm BH}=1,10), \\\\ \\\\ \\frac{40\\sqrt{2}}{7}\\cos^{4}{\\theta_{\\rm torus}}\\,\\,\\, (\\dot{m}_{\\rm BH}=10^{2}), \\\\ \\\\ \\frac{160}{3}\\cos^{9}{\\theta_{\\rm torus}}\\,\\,\\, (\\dot{m}_{\\rm BH}=10^{3}). \\end{array}\\right. \\end{equation} Here $r_{\\rm x}(\\geq r_{\\rm sub})$ is the distance from the X-ray source to the inner radius of the dusty torus. Assuming that the X-ray emission is isotropic, it is found that as $\\theta_{\\rm torus}$ and $\\dot{m}_{\\rm BH}$ increase, the X-ray emission from the corona begins to contribute to the IR emission from the dusty torus. However, we should keep in mind that the dust opacity for X-ray radiation is lower than that for optical-UV radiation (e.g., \\citealt{DL84,LD93}). Thus, the evaluated contribution of X-rays from the corona could be overestimated.  \\begin{figure} \\begin{center} \\includegraphics[height=5cm,clip]{f5.eps} \\end{center} \\caption{ Three typical models for the shape of the inner edge of the torus. $T_{\\rm d}$ and $T_{\\rm sub}$ are the dust temperature and sublimation temperature, respectively. Note that $r_{\\rm sub}(\\theta_{\\rm torus})$ is the sublimation radius at $\\theta=\\theta_{\\rm torus}$. The red thick lines and shaded portions denote the region of NIR emission (i.e., $T_{\\rm d}=T_{\\rm sub}$). The blue shaded portions show the region of no NIR emission (i.e., $T_{\\rm d} < T_{\\rm sub}$). } \\end{figure}  \\subsection{Physical models of the dusty torus} In $\\S 3.2$, we considered three typical models for the configuration of the inner edge of the dusty torus (see Fig. 5). We here discuss how these three types can be understood physically. The following two physical quantities may be closely related to the shape of the inner part of the dusty torus. One is the sublimation radius, $r_{\\rm sub}(\\theta)$, which determines the region where hot dust can survive. In order to maintain a geometrically thick structure, energy input from stars and/or supernovae in the dusty torus is necessary (e.g., \\citealt{WN02,Th05,KW08}). This energy input may occur outside of the critical radius, $r_{*}$, above which the torus is gravitationally unstable, if we adopt Toomre's stability criterion. Thus, the magnitude relation between $r_{\\rm sub}(\\theta_{\\rm torus})$ and $r_{*}$ may determine the structure of the inner part of the dusty torus, as in the following three cases. (i) If $r_{*} \\ll r_{\\rm sub}(\\theta_{\\rm torus})$, the inner radius of the torus, $r_{\\rm in}$, may be determined by $r_{\\rm sub}(\\theta)$. This corresponds to the case of Fig. 5(i). (ii) When $r_{*}\\approx  r_{\\rm sub}(\\theta_{\\rm torus})$, the structure of the dusty torus is like Fig. 5(ii). As mentioned above, the dust can survive at even $r < r_{\\rm sub} (\\theta_{\\rm torus})$, though the structure within $r_{*}$ is geometrically thin because there is no energy source in this region. Thus, the NIR emission from $r < r_{*}$ is negligible since the covering factor becomes small. That is, $r_{\\rm in}$ is determined by the radius for which the dusty torus is gravitationally unstable. (iii) If $r_{*}  \\gg r_{\\rm sub}(\\theta_{\\rm torus})$, the dust temperature of the whole inner edge of the dusty torus becomes lower than $T_{\\rm sub}$ because $r_{\\rm in}=r_{*}$ is much larger than $r_{\\rm sub}(\\theta)$. If this is the case, the structure of the dusty torus is similar to Fig. 5(iii). By exploring the relation between the region of nuclear starbursts and that of the hot dust, we may find evidence revealing the formation of the dusty torus.  \\subsection{Evolutionary tracks in $L_{\\rm IR,AGN}/L_{\\rm bol,disc} -\\theta_{\\rm torus}$ plane} We now discuss the evolutionary tracks of SMBH growth in the $L_{\\rm IR,AGN}/L_{\\rm bol,disc}$--$\\theta_{\\rm torus}$ diagram. As mentioned in $\\S 3.1$, although the formation mechanism of the obscuring torus is still a hotly debated issue (e.g., \\citealt{KB88,PK92,OU99,OU01,WN02,WU05, Wa09,Sc09}), \\citet{WN02} proposed a dusty torus supported by turbulent pressure, in which the turbulence is produced by SN explosions (see also \\citealt{KW08,Wa09}). In their model, the geometrical thickness of the torus, which is determined by the balance between the gravity of the central BH and the force due to turbulent pressure, is smaller for more massive BHs. This tendency is consistent with observations in which the covering factor of the dusty torus decreases with increasing BH mass (e.g., \\citealt{Ma07,No10}). This model also implies that $\\theta_{\\rm torus}$ increases with time as the BH grows. That is, the evolution proceeds from left to right in the $L_{\\rm IR,AGN}/L_{\\rm bol,disc}$--$\\theta_{\\rm torus}$ diagram (see Fig. 5). Here we consider two simple scenarios: (i) the super-Eddington growth dominated scenario where most of the mass of the SMBHs is supplied not through sub-Eddington accretion discs but through super-Eddington accretion discs. (ii) the sub-Eddington growth dominated scenario where the sub-Eddington accretion disc mainly feeds the SMBH. Figure 6 shows the evolutionary tracks for these two scenarios, assuming an initial opening angle of the torus of $\\theta_{\\rm torus,init}=45^{\\circ}$ and initially $\\dot{m}_{\\rm BH}=10^{3}$. We find from this figure that the AGNs stay in a super-Eddington phase ($\\dot{m}_{\\rm BH}>10^{2}$) for a long time (a wide range of $\\theta_{\\rm torus}$) in case (i), whereas in case (ii) the AGNs shift to a sub-Eddington phase when the torus is relatively thick. The two distinct scenarios can be clearly understood from Fig. 3 as follows: for case (i) the evolution proceeds from top left $\\to$ top right $\\to$ bottom right, whereas for case (ii) it goes from top left $\\to$ bottom left $\\to$ bottom right. That is, the AGNs never undergo a phase of extremely low $L_{\\rm IR,AGN}/L_{\\rm bol,disc}$ for case (ii). In contrast, in case (i), many AGNs can be identified as IR faint objects. The NIR luminosity of such AGNs is also very small. The dispersion of $L_{\\rm IR,AGN}/L_{\\rm bol,disc}$ as well as $L_{\\rm NIR,AGN}/L_{\\rm bol,disc}$ for case (i) is significantly larger than that for case (ii). If the majority of high-$z$ quasars ($z>7$) are NIR-faint, super-Eddington growth may be inevitable for SMBH growth in the early Universe. In order to compare the predictions with observations in detail, we must elucidate the evolution of $L_{\\rm NIR,AGN}/L_{\\rm bol,disc}$ based on the coevolution model of SMBH growth and a circumnuclear disc such as \\citet{KW08}. This is a subject for future work.  \\begin{figure} \\begin{center} \\includegraphics[height=6.5cm,clip]{f6.eps} \\end{center} \\caption{ Schematic diagram of two evolutionary scenarios, super-Eddington growth dominated scenario (case (i)) and sub-Eddington growth dominated scenario (case(ii)). The evolution proceeds from left to right, i.e., as the mass of the BH increases, $\\theta_{\\rm torus}$ increases (see text). We assume that the initial opening angle of the torus is $\\theta=45^{\\circ}$ and $\\dot{m}_{\\rm BH}=10^{3}$. } \\end{figure}  \\subsection{Comparison with observations} In this paper, we predicted that super-Eddington AGNs with $\\dot{m}_{\\rm BH}\\geq 10^{2}$ have relatively low $L_{\\rm NIR,AGN} /L_{\\rm bol,disc}$. So far, more than 30 quasars have been discovered at $z\\approx 6$ (e.g., \\citealt{Fa01,Fa06,Go06,Wi07,Wi10}). \\citet{Ji10} found that two are unusually NIR-faint quasars, i.e., their ratio of $L_{\\rm NIR}/L_{\\rm 5100}$ is one order of magnitude smaller than the average for ordinary quasars, although their properties are similar to those of low-$z$ quasars in the rest-frame ultraviolet/optical and X-ray bands (e.g., \\citealt{Fa04,Ji06,Sh06}). More interestingly, the two quasars have the smallest BH mass ($M_{\\rm BH} \\approx 10^{8}M_{\\odot}$) and the highest Eddington ratios ($L_{\\rm bol,AGN} /L_{\\rm Edd}\\sim 2$) of $z\\sim 6$ quasar samples. These features can be nicely explained by our interpretation, i.e., the faint NIR emission for high Eddington AGNs. We should note that since NIR-faint quasars can be explained by sub-Eddington AGNs with extremely thin tori (i.e., $\\theta_{\\rm torus}>85^{\\circ}$), we need to measure the thickness of a dusty torus to confirm whether NIR-faint quasars are really super-Eddington AGNs ($\\dot{m}_{\\rm BH} \\geq 10^{2}$). Based on the kinematics of cold molecular gas in a torus, it would be possible to evaluate the thickness of a torus by ALMA, because the ratio of velocity dispersion and rotational velocity indicates the scale height of the torus (see \\citealt{WN02,WT05}). On the other hand, \\citet{Ji10} interpreted these two objects as dust-free quasars. However, this seems to contradict the observations, which suggests that NIR-faint quasars possess super-metallicity as do other $z\\sim 6$ quasars and high-$z$ AGNs (e.g., \\citealt{Pe02,Fr03,Ju09,HF93,Na06a,Na06b,Ma09}). Also, \\citet{Ha10} recently reported quasars with unusually faint NIR emission at $z=1.5$--$3$, when the universe was 2--4 Gyr old. Thus, it seems likely that NIR-faint quasars are exclusively dust-free AGNs. As an alternative scenario, a large fraction of gas in a dusty torus would be ejected from host galaxies as a result of strong radiation pressure from the brightest AGNs, such as quasars \\citep{OU99,OU01,WU05}. In this case, NIR-faint quasars may be explained as being quasars without dusty tori around the SMBH. This possibility could be confirmed by investigating the presence of dusty tori in NIR-faint quasars using ALMA.  In the local universe, NLS1s are thought to be AGNs of a high $L_{\\rm bol, disc}/L_{\\rm Edd}$ system. One NLS1 (Ark 564) shows no NIR emission \\citep{RM06}, which is consistent with our predictions. However, most NLS1s tend to have NIR emission similar to or stronger than that of ordinary Seyfert galaxies (e.g. \\citealt{Ry07}). One of the reasons for the discrepancy between our predictions and the observations is that NLS1s with $\\dot{m}_{\\rm BH}\\geq 10^{2}$ may be absent (or their fraction is very small). However, some NLS1s have a high mass accretion rate with $\\dot{m}_{\\rm BH}\\geq 10^{2}$ (e.g., \\citealt{CK04,Ha08}). Another possibility is that these discrepancies may suggest misalignment between the accretion discs and the dusty torus, as mentioned in $\\S 3.1$. Lastly, we note the possibility of contamination of hosts and circumnuclear starbursts in the NIR band, if the host luminosity is comparable to or exceeds the AGN luminosity, as in Seyfert galaxies. Although it may be hard to find many NLS1s with low $L_{\\rm NIR,AGN}/L_{\\rm bol,disc}$ in current observations, in future studies it is worth examining this issue by NIR (rest frame) observations with high spatial resolution."
    },
    "1107/1107.0180_arXiv.txt": {
        "abstract": "Accreting millisecond pulsars show significant variability of their pulse profiles, especially at low accretion rates. On the other hand, their X-ray spectra are remarkably similar with not much variability over the course of the outbursts. For the first time, we have discovered that during the 2008 outburst of \\sax\\ a major pulse profile change was accompanied by a dramatic variation of the disc luminosity at almost constant total luminosity. We argue that this phenomenon is related to a change in the coupling between the neutron star magnetic field and the accretion disc. The varying size of the pulsar magnetosphere can influence the accretion curtain geometry and affect the shape and the size of the hotspots. Using this physical picture, we develop a self-consistent model that successfully describes simultaneously the pulse profile variation as well as the spectral transition. Our findings are particularly important for testing the theories of accretion onto magnetized neutron stars, better understanding of the accretion geometry as well as the physics of disc-magnetosphere coupling. The identification that varying hotspot size can lead to pulse profile changes has profound implications for determination of the neutron star masses and radii. ",
        "introduction": "\\sax\\ was first detected with the Wide Field Cameras on board the  {\\it BeppoSAX} satellite in 1996 \\citep{itz98}. In 1998, another outburst of \\sax\\ was observed with {\\it Rossi X-ray Timing Explorer} (\\xte) and the discovery of $\\approx 401$ Hz pulsations led to its identification as the first accreting millisecond pulsar \\citep[AMP;][]{WvdK98, CM98}. Since 1998, the source has been in outburst roughly every $2.5$ years (2000, 2002, 2005 and 2008). During the outbursts of \\sax\\ (and other AMPs), the magnetic field of the neutron star (NS) channels the accreted matter on to the stellar magnetic poles. The emitted radiation from these hotspots is then modulated at the stellar spin period, that results in coherent pulsations. In a typical outburst of \\sax\\ \\citep[see fig.2 in][]{HPC08} the source rises from quiescence in roughly five days, reaching a peak flux of about $2\\times10^{-9}\\, \\ergcm2s$ in the 2--25 keV range of the \\xte\\ Proportional Counter Array (PCA). After the peak, the source flux drops slowly (``slow decay\" stage) in the course of the next ten days, although the 2008 outburst had a double peaked light curve \\citep{HPC09}. The slow decay stage is followed by the ``rapid drop\" stage, where the flux drops by an order of magnitude in a couple of days after which the source exhibits several re-brightening episodes every five days or so (``flaring tail\" stage, see also \\citealt{PW09}, \\citealt{IP09}).  The system parameters of \\sax\\ are the best known among AMPs. The orbital period is $\\approx 2$ hours \\citep{CM98, BRD09, HPC09} and the companion star has a very low mass of $M_{\\rm c} \\sim 0.05 \\Msun$ \\citep{BC01}. The analysis of type I X-ray bursts by \\citet{GC06} gave a distance estimate of $3.5 \\pm 0.1$ kpc. The inclination has been constrained to $i=$36$^\\circ$--67$^\\circ$ using optical observations \\citep{DHT08} and the X-ray analysis \\citep{IP09} of the 2002 outburst led to a similar constraint of $i \\approx 50^\\circ$--$70^\\circ$. By studying the long term timing evolution of the system, \\citet{HPC08} constrained the neutron star surface magnetic dipole field to a range of $\\Bs = (0.4$--$1.5) \\times 10^8\\,{\\rm G}$. Another estimate of $ \\Bs = (0.8 \\pm 0.5) \\times 10^8 k_{\\rm A}^{-7/4}\\,{\\rm G}$ from pulse profile variability, was derived by \\citet{IP09}, where the factor $k_{\\rm A} \\approx 0.5$ \\citep{long05}. Also, \\citet{BDM06} gave an estimate of $ \\Bs \\sim (3.5 \\pm 0.5) \\times 10^8\\,{\\rm G}$ and \\citet{PWvdK09} obtained $\\Bs = (2$--$3) \\times 10^8\\,{\\rm G}$.   The energy spectra of \\sax\\ has been studied extensively \\citep[e.g.][]{GR98,GDB02,PG03,IP09}. It was pointed out already by \\citet{GR98} that the spectral shape remains remarkably similar when flux changes by more than an order of magnitude throughout the outburst. The energy spectrum is roughly flat (photon index $\\Gamma \\approx 2$) with a cutoff/roll-over at energies above $\\sim 50$ keV. Such spectra -- also seen in other AMPs \\citep{P06} -- are most likely produced by thermal Comptonization of soft seed photons originating from the heated stellar surface by hot electrons in the accretion shock \\citep{GDB02,PG03,IP09}. Below $\\sim 5$ keV an additional $T \\approx 0.5$ keV thermal component is visible that is associated with the heated surface of the NS. \\xmm\\ observations have revealed another {\\it non-pulsating} cooler thermal component due to the accretion disc \\citep{PRA09}. The X-ray emission from the NS surface irradiates the accretion disc, that causes spectral hardening above 10 keV due to Compton reflection \\citep{IP09}. This irradiation also produces an iron line at 6.4 keV and the modelling of the line profiles from \\xmm\\ and {\\it Suzaku} observations by \\citet{PDD09} and \\citet{CAP09} has given tight constraints on inclination and the accretion disc truncation radius. We note, however, that these results depend strongly on the continuum spectral model and data reduction issues (see \\citealt{NDC10}, for discussion).   Although many aspects of AMP physics are known, one of the unresolved puzzles is the origin of sudden pulse profile changes in \\sax\\ \\citep{HPC08,HPC09} and other AMPs. There are many mechanisms that can cause such changes \\citep{P08} and they contribute to the ``timing noise\" seen in many AMPs \\citep[e.g.][]{BDM06,pap07,RDB08,PWvdK09, PIA09}. One of these remarkable pulse profile changes occurred in the 2008 outburst of \\sax\\ \\citep{HPC09}. In the beginning of the slow decay stage, the pulse profile was rather symmetric showing only small harmonic content. However, on September 27 the fundamental pulse amplitude decreased by almost 50 per cent, while no change was seen in other observed quantities, especially the observed flux and phases of the harmonics. A few days later on October 2 (MJD 54742), the pulse profile quickly morphed to have two peaks. This time the pulse profile transition was accompanied by a decay in the observed flux and jumps in pulse phases. The double peaked profile was seen only for about three days, because on October 6 (MJD 54746), right before the onset of the rapid drop stage, the fundamental amplitude increased significantly and the profile changed again to a shape similar to what was seen before September 27. This behaviour cannot be explained by a change in the accretion disc truncation radius that was suggested by \\citet{IP09} and \\citet{PIA09}. Although this model explains the jump in the fundamental phase (and the pulse profile change), that is associated with the rapid drop stage of the 2002 outburst of \\sax\\ \\citep{BDM06}, it cannot be the case here simply because the observed flux (which can be related to the truncation radius) in these sudden pulse profile changes is almost an order of magnitude higher and it remains nearly constant \\citep[see fig. 1 in][]{HPC09}. Among the 13 known AMPs \\citep[see][and references therein]{Pat10}, some sources show similar jumps in the fundamental pulse amplitudes as in the 2008 outburst of \\sax. A few examples include XTE J0929--314 \\citep{GC02}, the first eclipsing AMP SWIFT J1749.4--2807 \\citep{MS10,ACP11} and SWIFT J1756.9--2508 \\citep{PAM10}. Also, similar jumps in the fundamental pulse amplitude was seen in the 2002 and 2005 outbursts of \\sax\\ (see fig. 6 in \\citealt{IP09} and fig. 1 in \\citealt{HPC08}).  The variability in pulse amplitudes and profiles are sometimes accompanied with jumps in pulse phases \\citep{pap07,RDB08,PWvdK09}. In most cases, this timing noise is related with changes in the observed flux \\citep{BDM06, IP09, PWvdK09} and the most notable changes usually occur in the fundamental phase \\citep{BDM06,pap07,RDB08,PWvdK09}. This type of timing noise is most likely caused by mass accretion rate dependent change of the hotspot location on the NS surface \\citep{LBW09}, which explains the overall trends with flux and pulse phase residuals in \\sax\\ \\citep{HPC08} and in many other AMPs \\citep{PWvdK09}. However, although the general trends can be explained \\citep{PWvdK09}, there are several AMPs where additional timing noise is clearly present. Most notable cases are the 2002, 2005 and 2008 outbursts of \\sax, XTE J1807--294 and XTE J0929--314 (see fig. 3 in \\citealt{PWvdK09} and fig. 1 in \\citealt{HPC09}). In these cases another mechanism -- that is not related to changes in mass accretion rate -- must be responsible for the timing noise.  In this paper, we present our spectral and timing analysis of the 2008 outburst of \\sax\\ using \\swift\\ and \\xte\\ data. We find evidence that the sudden changes in the pulse profiles and in pulse amplitudes of \\sax\\ are driven by a changing interaction between the NS magnetic field and the accretion disc. ",
        "conclusions": "\\subsection{Implications for spin-up torques during the outburst}  The changes in the interaction between the disc and the magnetosphere can have a large effect on the torque that spins up (or spins down) the NS. Several different models have been developed to compute the torques \\citep[e.g.][]{GL79a, GL79b, W87, LRB95, RFS04, MP05, KR07, MP10}. Generally, in addition to the ``accretion torque\" $\\tau_{\\rm a} = \\dot{M}\\sqrt{G \\Mns \\Rt}$ \\citep[e.g.][]{MP10}, the disc--magnetosphere  coupling can cause an additional ``magnetic torque,\" which is strongly dependent on the nature of the interaction and varies significantly between the models. Nevertheless, it can be concluded that small variations of $\\Rm$  independent of the mass accretion rate (as observed here for \\sax) can cause the NS spin-up or spin-down depending on the geometry and position of $\\Rm$ relative to $\\Rco$.  Also the data suggests, that only those models that predict that $\\omega \\approx \\omegans$ in the magnetosphere can be valid to explain the observed spectral transition. So observations of AMPs -- as presented here -- can be used to test different hypothesis and constrain the nature of the disc--magnetosphere coupling.  \\subsection{Implications for the timing noise in AMPs}  Although the pulse amplitude dropped significantly, the pulse phase remained constant during the transition. This is actually not surprising, based on the fact that the magnetic inclination seems to be about $\\theta \\approx 10\\degr$. As the hotspot shape is most likely close to a ring \\citep{RU04}, the change in $\\Rm$ should only affect the location of the inner edge of the hotspot $\\rhoin$, but not change the location of the hotspot on the NS surface. However, the situation might be different for other AMPs, where $\\theta$ can be larger. For such AMPs, the hotspot shapes are most certainly not symmetric \\citep{RU04}, and changes in $\\Rm$ might result in a change in the shape, size, latitude and --most importantly-- the hotspot longitude. Shifts in the hotspot longitude can especially cause large jumps in the observed pulse phases \\citep[see][for discussion]{LBW09}. We speculate that such $\\Rm$--dependent motion of the hotspots (and the associated phase jumps) could be the origin of at least part of the timing noise in AMPs. Observationally, such mechanism could cause the outliers in the X-ray flux -- phase residual relations in XTE J1807--294 \\citep{RDB08,PWvdK09,PHW10}, XTE J0929--314 \\citep{GC02,PWvdK09} and IGR J17511--3057 (\\citealt{RPB11}; \\citealt{IKP11}), whereas the overall X-ray flux -- phase residual trends can be caused by $\\Mdot$--dependent motion of the hotspots \\citep{LBW09,PWvdK09,PHW10}. Also, we stress that an additional contribution to the timing noise can be caused by $\\Mdot$--dependent variation of the truncation radius $\\Rt$, as it affects the visibility of the secondary spot \\citep{IP09,PIA09}. These factors do not however exclude the possibility that some AMPs do spin-up during X-ray outbursts (as is most likely the case with IGR J00291+5934, e.g. \\citealt{FKP05,BDL07,Pat10b,HGC11,PRB11}).  \\subsection{Possible caveats} \\label{sec:caveats}  \\begin{enumerate} \\item The truncation radius of $\\Rt \\approx 20$ km was obtained from pulse profile modelling. We assumed that this radius is not affected by the change in the disc--magnetosphere coupling that we propose to be the reason for the transition. We also assumed that the outer edge of the ring-shaped hotspot $\\rhout$ is determined from equation (\\ref{xiout}), which is not necessarily the case as the field topology should differ from a pure dipole. These factors can change the constraints on the other parameters. However, a detailed analysis of these factors should be done by 3D MHD simulations \\citep[as in][]{RU04}, but such a study is clearly beyond the scope of the present paper.  \\item As the Ohmic dissipation in the magnetosphere is small, we proposed that $\\Rdisc \\sim \\Rm$. However, there are some uncertainties in deriving $\\Rdisc$ in addition to those related to spectral modelling in Section \\ref{modelling}. The main systematic uncertainty in estimating $\\Rm$ from $\\Rdisc$ comes from the fact that $\\Tdisc$ is in reality a colour temperature and not the effective temperature of the disc. Also, the inner disc boundary condition might differ significantly from what is assumed in the {\\sc diskbb} model, where the radial dependence of temperature is $T(r) \\propto r^{-3/4}$. These effects could be, in principle, accounted for to obtain the corrected inner disc radius \\be R'_{\\rm disc} =  \\xi f_{\\rm col}^2 \\Rdisc, \\ee where  $\\xi$ accounts for the fact that the modelled inner disc temperature $\\Tdisc$ does not actually occur at $\\Rdisc$ and $f_{\\rm col}$ is the colour correction factor \\citep{G99}. The colour correction factor was computed to be $f_{\\rm col} \\approx 1.7$ by \\citet{ST95} for black hole systems in the soft state and it seems to be weakly dependent of the mass accretion rate \\citep[see][]{DBH05}. The correction factor $\\xi = 0.37$ was computed by \\citet{G99} by assuming a zero torque boundary condition at the inner edge of the disc in the pseudo-Newtonian potential. However, because of the different inner disc boundary conditions for accretion onto an AMP $\\xi$ can be much different especially in the beginning of the outburst, where the spectral modelling indicates that $\\Rm \\gtrsim \\Rco$ and therefore $\\omega_{\\rm K} < \\omegans$. In this case, there should be a region around $\\Rm$ where a large jump from the Keplerian rotation to the co-rotation should occur (LRB95, \\citealt{LRN10}). An estimate of the size of this ``transition region\" $\\Delta r$ was given by LRB95 \\be \\Delta r \\sim \\gamma_{\\rm c} \\frac{\\eta_{\\rm t} c_{\\rm s}}{\\omega_{\\rm K} \\nu_{\\rm t}}. \\ee For a thin accretion disc $c_{\\rm s}  / v_{\\rm K}  \\sim H / r \\ll 1$, so at $r=\\Rm$ the transition region $\\Delta r \\sim H \\ll \\Rm$ assuming that $\\eta_{\\rm t}  \\sim \\nu_{\\rm t}$ and $\\gamma_{\\rm c}  \\sim 1$. Therefore, we speculate that its effect on the observed radiation flux is small due to the small size, but the exact value of the correction factor $\\xi$ is uncertain. Furthermore, as the accretion disc is irradiated by the emission of the hotspot and part of the hard emission is absorbed by the disc, the observed colour temperature can be altered and therefore affect $\\Rdisc$. But because the amplitude of Compton reflection $\\refl$ is poorly constrained, it is not possible to accurately take this effect into account. All these unknown factors cause a systematic error in $\\Rdisc$, which makes an accurate estimation of $\\Rm$ currently infeasible.  Similarly, the hotspot temperature $\\Tbb$ is also a colour temperature (at the infinity), but $f_{\\rm col}$ is most likely the order of unity in this case \\citep[see][and references therein]{GP05} and the actual hotspot size can be larger depending on the stellar compactness and the hotspot geometry \\citep[e.g.][]{PG03,GP05,IP09}. These considerations only change the reported results quantitatively, but the qualitative result of the hotspot size variation shown in Fig. \\ref{fig:parameters} is not affected.  \\item In our modelling we did not consider any complicated spot shapes. Based on the simulations of \\citet{RU04} a crescent shaped spot might have been more appropriate to use. However, this would have required several new parameters to the model, which is not justified based on the quality of the data. This cannot be improved because the pulse profile changes so rapidly that co-adding more data would not be appropriate. Also by selecting a broader range in energies than the 9.8--23.2 keV band to improve the statistics cannot be done easily, because the blackbody component (with a different emission pattern) will start contributing from below and there are not enough photons above this range. Therefore, the only way to address this issue is by making extensive simulations and to extract pulse profiles from all the outbursts of \\sax\\ and then simultaneously fit them. \\end{enumerate}"
    },
    "1107/1107.5064_arXiv.txt": {
        "abstract": "We consider the astrophysical constraints on the gravitational-wave driven r-mode instability in accreting neutron stars in low-mass X-ray binaries. We use recent results on superfluid and superconducting properties to infer the core temperature in these neutron stars and show the diversity of the observed population. Simple theoretical models indicate that many of these systems reside inside the r-mode instability region. However, this is in clear disagreement with expectations, especially for the systems containing the most rapidly rotating neutron stars. The inconsistency highlights the need to re-evaluate our understanding of the many areas of physics relevant to the r-mode instability. We summarize the current status of our understanding, and we discuss directions for future research which could resolve this dilemma. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5314_arXiv.txt": {
        "abstract": "We investigate in detail the probability distribution function (pdf) of the proper-motion measurement errors in the SDSS+USNO-B proper-motion catalog of \\citet{mun04} using clean quasar samples. The pdf of the errors is well-represented by a Gaussian core with extended wings, plus a very small fraction ($<0.1\\%$) of  ``outliers''. We find while formally the pdf could be well-fit by a five-parameter fitting function, for many purposes it is also adequately to represent the pdf with a one-parameter approximation to this function. We apply this pdf to the calculation of the confidence intervals on the true proper motion for a SDSS+USNO-B proper motion measurement, and discuss several scientific applications of the SDSS proper motion catalogue. Our results have various applications in studies of the galactic structure and stellar kinematics. Specifically, they are crucial for searching hyper-velocity stars in the Galaxy. ",
        "introduction": "\\label{sec:introduction}  As of the eighth Data Release (hereinafter DRn, where n is the release number) of the Sloan Digital Sky Survey (hereinafter SDSS) the survey has released imaging for 14,555 deg$^2$, or over a third of the sky \\citep{gun98,yor00,lup01,sto02,aih11}. The imaging catalog contains almost half a billion distinct detected objects down to a $50\\%$ completeness limit of $r=22.5$ for point sources, and the survey's 1.8 million catalogued spectra provide several well defined samples of galaxies, quasars, stars, and other objects \\citep{aih11}.  Along with the photometric and spectroscopic data, SDSS also provides proper-motion (PM) data \\citep[hereinafter Munn et al. catalog]{mun04}, produced by matching the SDSS point source detections with earlier observations, including the USNO reductions of the  Palomar Observatory Sky Surveys (POSS-I and POSS-II), which span about 50 yr in time \\citep{monet03}. In this catalog, the USNO proper-motion system is re-calibrated and made absolute using SDSS galaxies, and these proper motions (called here SDSS+USNO-B PM) are computed including both SDSS and USNO-B positions. The resultant catalog is $90\\%$ complete to $g<19.7$, and  has a less than $0.5\\%$ contamination rate. The systematic errors are on the order of 0.1 \\mpy, and the statistical errors are roughly 3-4 \\mpy ~in each component of the PM. \\cite{mun04} compared their results with those of the revised New Luyten Two-Tenths catalog (rNLTT; \\citealt{gou03,sal03}). \\cite{bon10} carried out a further comparison of the SDSS+USBO-B proper motions with those of a sample of stars in the North Galactic Pole region \\citep{maj92} and with proper motion measurements made using data from SDSS Stripe 82 \\citep{bra08}. These independent measurements are expected to have different systematic errors from the SDSS+USNO-B PM. The results show that the median differences and the rms scatter of these comparisons agree with expectation and that the SDSS+USNO-B PM measurements are reliable at roughly the stated errors.  Proper-motion measurement at this level of accuracy are useful in several ways. Samples of objects can be defined using the reduced proper motion diagram to separate classes of objects with intrinsically similar colors, spectra and apparent magnitudes but very different proper motion distributions due, for example, to different luminosities or different kinematics. This is used in several target selection algorithms in the SDSS projects, most notably the Sloan Extension for Galactic Understanding and Evolution \\citep[SEGUE I and II]{yan09}, both to help find very nearby objects (with high PM) (e.g. \\citealt{lep08}) and distant giants (with low PM). A variation of this method has been explored by defining a ``reduced-proper-motion'' (for example in the $r$-band the reduced-proper-motion would be defined as $r_{\\rm RPM}=r+5\\log{\\rm PM}$, \\citealt{sal03}), while the reduced proper motion diagrams are not as useful at faint magnitudes probed by SDSS as they are at the (traditionally used) bright end \\citep{ses08}. Most importantly, PM, along with even crude distance determinations for stars in the Galaxy, provide two-dimensional velocity information; if in addition radial velocity measurements are available, one has full three-dimensional velocity information. A one-sigma accuracy of 3 \\mpy corresponds to a transverse velocity accuracy of about 15 km/s at 1 kpc and 150 km/s at 10 kpc. Thus the PM measurements are {\\it statistically} very useful for studying the kinematics of the thick disk and halo, for which the velocity dispersions are of this order at these distances---provided one understands the errors well. Recent years have witnessed an increased interest in studies of stellar kinematics and Galactic structure, spurred both by the increasing sophistication of galaxy formation simulations and the availability of the large photometric sample of stars in SDSS together with a significant subsample with radial velocities and chemical composition information obtained as part of its SEGUE subsurveys. With the SDSS+USNO-B PM measurement and radial velocities, \\citet{bon10} selected a sample of main-sequence stars with $r<20$, while  \\citet{car10} selected a subsample of SDSS calibration stars (mostly metal-poor turnoff F stars) to study the kinematics of both the galactic thin and thick disks and the halo (see also \\citealt{smi09,sch09,fuc09}). Many more studies of this sort are underway using the SEGUE II data.  To do this job really well clearly requires detailed knowledge of the distribution of the proper motion, distance and radial velocity errors, since for most of these applications the proper-motion related velocity errors are of the same order as the velocities themselves. Furthermore, in most cases the error in the tangential velocity due to the proper motion and distance error is the dominant contributor to the total velocity error. In addition, if one is interested (and one usually is) in extreme velocities, understanding the behavior of the distribution in the tails of the pdf is crucial. We approach this problem in the way others have done in the past, but with the goal from the outset of understanding the pdf in as much detail as the sample sizes and systematics will allow.  Quasars are sufficiently distant that they should not have any measurable PM in SDSS. Thus, the measured PM of quasars are just the PM measurement error. By studying the PM of a large sample of quasars, we can probe the statistical properties of the PM errors, and study the dependence on various observational and instrumental parameters. Although the spectral energy distributions for quasars differ from those of the point sources for which PM is of interest (stars), the major contributors to the systematic error in the PM measurement, including: the difficulty in centroiding sources on photographic plates; errors due to unresolved or partially resolved projected nearby objects; errors in the primary reference catalogs; systematic errors in the UCAC catalog  \\citep{zacharias10};  and charge transfer effects in the SDSS astrometric and photometric detectors \\citep{mun04}, work in essentially the same way for all point sources. Thus the error distribution for all point sources will be very similar (\\citealt{bon10}, but caution is still needed when applying the results based on quasars to stars, see Section \\ref{sec:fitting} and \\ref{sec:summary} for details). Along with the publication of the PM catalog, \\citet{mun04} used spectroscopic quasars in the SDSS DR1 \\citep{schneider02} to study the mean and variance of the SDSS+USNO-B PM error, but due to the limited sample size  were not able to study the dependence on magnitude. With a much larger spectroscopic quasar sample in SDSS DR7, \\citet{bon10} addressed this issue again and presented the width $\\sigma$ in the Gaussian distribution as a function of $r$-band magnitude. They assumed the error distribution was Gaussian and did not investigate the form of the distribution function; in particular, they did not try to characterize the wings of the distribution.  In this work, we analyze the full error distribution of the SDSS+USNO-B PM (with the corrections presented by \\citealt{mun08}, which corrected an error in the calculation of the proper motions in right ascension). We do this using clean quasar samples which we define in such a way that corresponding clean stellar samples are easily constructed. We fit the PM error distribution in the entire magnitude range by a Gaussian distribution for the core plus a wing function, and derive the dependence of the function parameters on magnitude. Using this fitting function, we calculate the significance of a given PM measurement. We also quantify (or at least place upper limits on) the fraction of the outliers in the PM measurement which survive the cleaning process and for one reason or another clearly do not belong to the main distribution. We then discuss various issues which involve applying the analysis of the PM error distribution to other samples.  The structure of this paper is as follows. In Section~\\ref{sec:sample} we introduce the quasar samples and the criteria for selecting reliable PM. We then fit the PM error distribution by parametrized analytic functions of both complex and and simplified forms in Section~\\ref{sec:fitting}, and calculate the significance of measured PM in Section~\\ref{sec:significance}. We summarize our results and discuss the issues related to using them in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have investigated the proper-motion measurement errors in the SDSS+USNO-B proper-motion catalog \\citep{mun04} by analyzing the proper-motion distributions of several recent SDSS quasar samples \\citep{ric09,bon10,sch10,bov11}. The sample defined by \\citet{sch10}was determined to be the cleanest sample and has the largest sample size. We bin the data into six $g$ magnitude (not extinction corrected) bins, and fit analytic functions for the PM distribution in each bin. We find that, while a six-parameter fitting formula (Equation~\\ref{eq:f6}) describes the quasar PM distribution well, a simpler (normalized) function with one free parameter (Equation~\\ref{eq:f2}) also gives reasonably good results. Based on this fitting function for the PM error distribution, we calculate the probability that an object with a measured PM has a true PM $p_{\\rm true}$ larger than a given threshold $p_{\\rm true}^\\prime$. Cutting the probability $F(p_{\\rm true}>p_{\\rm true}^\\prime)$ at several confidence levels, we calculate the ``most likely PM'' for a given measured PM.  The analysis raised several issues which we would like to stress:  \\begin{enumerate} \\item Our PM error analysis can only be applied to {\\it clean} PM subsamples from the SDSS+USNO-B catalog, which satisfies our strict clean PM criteria ({\\tt match = 1}, {\\tt sigRa < 525 \\&\\& sigDec < 525}, {\\tt dist22 > 7}, and most importantly, {\\tt nFit = 6}, which requires that the object was detected on all five USNO-B plates). Experiments show that weakening this set of conditions (specifically, the original criteria in \\citet{mun04}: {\\tt match = 1} and {\\tt sigRa < 350 \\&\\& sigDec < 350}) generally results in a broader Gaussian core width (up to $10\\%$), a worse $\\chi^2$ statistics in the fitting, and a significantly higher fraction of the outliers (by an order of magnitude). Studies which are sensitive to the tail of the PM error distribution and the fraction of outliers, such as searching for extremely high velocity stars in the Galaxy, will be severely affected if using this weaker set of conditions instead of ours.  \\item The PM error distribution derived here applies to $g$ magnitudes (without extinction correction) in the range $\\sim17-20.5$. Specifically, the fitting functions Equation~\\ref{eq:f2} and \\ref{eq:sigma-g} are only good for $17.5\\leq g\\leq20.5$. Though using a constant $\\sigma$ for brighter objects is probably appropriate, using a constant or extrapolating our results for fainter magnitudes is not (Section~\\ref{sec:fitting}).  \\item While we have derived a fitting function for the entire PM error range, we note that there is a small fraction of ``outliers'' (defined here to be PM error $>30$ \\mpy), which apparently do not belong to the derived error distribution. Even for our cleanest quasar sample, which has passed spectroscopic visual inspection \\citep{sch10}, and with our very conservative PM culling conditions, there is still a small number of outliers, up to $\\sim0.1\\%$, in the faintest magnitude bands. We do not speculate on the origin of these, but they are likely present in any sample chosen by our criteria.  \\item While we fit the PM error distribution as a function of magnitude, the readers should bear in mind that the PM error may depend on other parameters as well, such as color and position of the objects on they sky. The PM error distribution in this work is best applicable to objects in the color range similar to the quasar sample we use to get the distribution (i.e., blue stars, Fig. \\ref{fig:color}), though our result could be considered as a conservative limit for redder objects. While the median and rms for the PM error only weakly depend on the position on the sky, we do not have a large enough sample to investigate the dependence of the distribution on the position on the sky. The distributions of the PM error components in the longitudinal and latitudinal are very similar to each other (Fig. \\ref{fig:pmlpmb-histo}), and the correlation between the two is negligible compared to the total random and systematic errors.  \\item Last, we note that the PM error estimate provided by \\citet[{\\tt pmraerr} and {\\tt pmdecerr}]{mun04} is in rough agreement with the $\\sigma$ fitted from the quasar PM distribution (Fig. \\ref{fig:sigma-g}). The agreement is within $20\\%$, and better at the faint end. On the other hand, the major drawback of the catalog-provided error estimate is that it does not address the non-Gaussian tail. We recommend that future analysis using Munn et al. proper motions, especially the ones for which the tail of the error distribution is important, should use this work to estimate the error instead of using the catalog-provided error estimate. \\end{enumerate}  This investigation of the PM error are useful in many ways. For example, we are studying the kinematics of stars in SEGUE II, focusing on these hyper velocity ones with velocity exceeding the escape velocity of the Galactic potential. In this case, it is crucial to have an accurate PM error distribution to judge the significance of the hyper velocity for these escapers."
    },
    "1107/1107.3532_arXiv.txt": {
        "abstract": "In this article we investigate the properties of the FLRW flat cosmological models in which the cosmic expansion of the Universe is affected by a dilaton dark energy (Liouville scenario). In particular, we perform a detailed study of these models in the light of the latest cosmological data, which serves to illustrate the phenomenological viability of the new dark energy paradigm as a serious alternative to the traditional scalar field approaches. By performing a joint likelihood analysis of the recent supernovae type~Ia data (SNIa), the differential ages of passively evolving galaxies, and the Baryonic Acoustic Oscillations (BAOs) traced by the Sloan Digital Sky Survey (SDSS), we put tight constraints on the main cosmological parameters. Furthermore, we study the linear matter fluctuation field of the above Liouville cosmological models. In this framework, we compare the observed growth rate of clustering measured with those predicted by the current Liouville models. Performing a $\\chi^{2}$ statistical test we show that the Liouville cosmological model provides growth rates that match sufficiently well with the observed growth rate. To further test the viability of the models under study, we use the Press-Schechter formalism to derive their expected redshift distribution of cluster-size halos that will be provided by future X-ray and Sunyaev-Zeldovich cluster surveys. We find that the Hubble flow differences between the Liouville and the $\\Lambda$CDM models provide a significantly different halo redshift distribution, suggesting that the models can be observationally distinguished. ",
        "introduction": "\\label{sec:intro} The comprehensive study carried out in recent years by the cosmologists has converged towards a cosmic expansion history that involves a spatially flat geometry and a recent accelerating expansion of the Universe (see \\cite{Teg04,Spergel07,essence,Kowal08,Hic09,komatsu08,LJC09,BasPli10} and references therein). From a theoretical point of view, an easy way to explain this expansion is to consider an additional energy component, usually called dark energy (hereafter 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). Indeed the so-called spatially flat concordance $\\Lambda$CDM model, which includes cold dark matter (CDM) and a cosmological constant ($\\Lambda$), fits accurately the current observational data and thus it is an excellent candidate to be the model which describes the observed Universe.  Nevertheless, the identification of $\\Lambda$ with the quantum vacuum has brought another problem: the estimate of theoretical physicists that the vacuum energy density should be 120 orders of magnitude bigger than the measured $\\Lambda$ value. This is the ``old\" cosmological constant problem (CCP) \\cite{Weinberg89}. The ``new\" problem \\cite{Peebles03} asks why is the vacuum density so similar to the matter density just now? Many solutions to both problems have been proposed in the literature \\cite{Zlatev99,AASW09,Egan08}. An easy way to overpass the above theoretical problems is to replace the constant vacuum energy with a DE that evolves with time. Nowadays, the physics of DE is considered one of the most fundamental and challenging problems on the interface uniting Astronomy, Cosmology and Particle Physics and indeed in the last decade there have been theoretical debates among the cosmologists regarding the nature of this exotic component.  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 resemble the DE~\\cite{Peebles03,Pad03,Jassal,SR,Xin,SVJ}. This class of DE models have been widely used in the literature due to their simplicity. Notice that DE models with a canonical kinetic term have a dark energy EoS parameter $-1\\leq w_{\\phi}<-1/3$. Models with ($w_{\\phi}<-1$), sometimes called phantom DE~\\cite{phantom}, are endowed with a very exotic nature, like a scalar field with negative kinetic energy. 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, 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_{\\Lambda}=\\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}\\sim10^{19}$ GeV, the corresponding mass of $\\phi$ is expected in the ballpark of $m_{\\phi}\\sim\\,H_{0}\\sim10^{-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,Basi09}; 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}\\sim10^{-33}~{\\rm eV}$.  Despite the above difficulties there is a class of viable models of quintessence usually called dilatonic that are based on supersymmetry, supergravity and string-theory and which can protect, for specific potentials, the light mass of quintessence (for a review see Ref.~\\cite{Ame10} and references therein). In particular, in string theory, gauge and gravitational couplings are connected with the vacuum expectation value of scalar field called dilaton $\\phi$~\\cite{Gas}. In this context, it has been proposed by some of us a specific model for the DE. This DE model being associated with a rolling dilaton field that is a remnant of this non-equilibrium phase described by a generic non-critical (Liouville) string theory~\\cite{emnw,diamandis,diamandis2,lmn}. We call this scenario either $Q$-cosmology or Liouville cosmology. This pattern is based on non-critical (Liouville) strings~\\cite{aben,ddk,emn}, which offer a mathematically consistent way of incorporating time-dependent backgrounds in string theory. We note in passing that the presence of time dependent dilaton fields at late eras of the Universe, that characterizes $Q$-cosmology, may lead to phenomenologically interesting extensions of minimal supergravity models. The latter predict less dark-matter relic abundances than the conventional $\\Lambda$CDM model, thereby allowing for more room for supersymmetry in collider (such as LHC) tests of the models~\\cite{lmn,lahanas2010}.   In this article we investigate the main dynamical properties of the $Q$-cosmological model from the point of view of current astrophysical constraints, extending non-trivially earlier analyses~\\cite{mitsou,mitsou2}. Initially, a joint statistical analysis, involving the latest observational data [SNIa~\\cite{Kowal08}, $H(z)$~\\cite{Stern10} and BAO~\\cite{Eis05,Perc10,Kazin10}] is implemented. Secondly, we attempt to compare the matter fluctuation field of the $Q$-cosmology with that of the $\\Lambda$CDM models by computing the growth rate of clustering. Finally, by using the Press-Schechter formalism~\\cite{press} and recently derived functional forms for the dark matter halo mass function~\\cite{Reed}, we show that the evolution of the cluster-size halo abundances is a potential discriminator between the $Q$-cosmology and $\\Lambda$CDM models. We would like to stress here that the abundance of collapsed structures, as a function of mass and redshift, can be accessed through observations~\\cite{Evra}. Indeed, the mass function of galaxy clusters has been measured based on X-ray surveys~\\cite{Borg01, Reip02, Vik09}, via weak and strong lensing studies~\\cite{Bat98, Dahle06, Corl09}, using optical surveys, like the SDSS~\\cite{Bah03, Wen10}, as well as, through Sunayev-Zeldovich (SZ) effect~\\cite{Taub05}. In the last decade many authors have found that the abundance of collapsed structures is affected by the presence of a dark energy component~\\cite{Wein03,Liberato,manera,Abramo07,Fran08,Sch09,Mort09,Rap10,Pace10,Alam10,Khed10,BPL10,Lomb10}.  The structure of our paper is as follows. The basic elements of the Liouville cosmological model are presented in Section~\\ref{sec:theo}, where we also introduce the cosmological equations for a flat Friedmann-Lemaitre-Robertson-Walker (FLRW) geometry. In Section~\\ref{sec:likelihood}, a joint statistical analysis based on SNIa, $H(z)$ and BAO is used to constraint the Liouville model free parameter. The issue related with the effective DE equation of state parameter and thus with the cosmic expansion is presented in Section~\\ref{sec:expa}. The linear growth factor of matter perturbations is discussed in Section~\\ref{sec:growth}, while in Section~\\ref{sec:cluster}, we discuss and compare the corresponding theoretical predictions regarding the evolution of the cluster abundances. Finally, the main conclusions are summarized in  Section~\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu} In this paper, we have studied the overall dynamics of the Liouville cosmological model in which the dark energy component is associated with the dilaton. In this framework, we first performed a joint likelihood analysis in order to put tight constraints on the main cosmological parameters by using the current observational data [SNIa, BAOs and $H(z)$]. We found that the above models can accommodate a late-time accelerated expansion.  Secondly, we performed an additional detailed statistical analysis based on the observed growth rate of clustering and found that the predicted growth rate of various Liouville cosmological models match well the observed growth rate.  The redshift-dependence of the halo abundances, predicted by the Liouville model, is significantly different from that of the flat $\\Lambda$CDM cosmology, at relatively high redshifts ($z\\gtrsim 0.6$). Future cluster of galaxy surveys can therefore be used to discriminate the Liouville model from other contender DE models."
    },
    "1107/1107.3829_arXiv.txt": {
        "abstract": " ",
        "introduction": "Extra-dimensional theory is a fascinating approach that affords a compelling candidate solution to the hierarchy problem~\\cite{ArkaniHamed:1998rs}. Indeed, the discovery of an evidence of extra-dimension as well as the Higgs particle and supersymmetry (SUSY) is one of the important missions of the CERN Large Hadron Collider (LHC) experiment. In the extra-dimensional theory, there are also some alternatives to the ordinary electroweak symmetry breaking (EWSB) mechanism in the standard model (SM), such as the gauge-Higgs unification (GHU)~\\cite{Fairlie:1979at}, the little Higgs~\\cite{ArkaniHamed:2001nc}, the Higgsless~\\cite{Csaki:2003dt}, and the Dirichlet Higgs~\\cite{Haba:2009pb} models, and so on. Furthermore, extra-dimensional models can give phenomenologically interesting features and predictions for physics beyond the SM, for example, the presences of Kaluza-Klein particles and a candidate for dark matter (DM)~\\cite{Zwicky:1933gu} from the Universal Extra-Dimensions (UED) model~\\cite{Antoniadis:1990ew}, and deviations of couplings of the Higgs~\\cite{Haba:2009pb,Hosotani:2008tx}.  The elucidation of the origin of DM, which governs about 23\\% of the Universe~\\cite{Komatsu:2010fb}, is one of important goals of the particle physics today. A large number of DM candidates have been discussed in literature.\\footnote{See ref.~\\cite{Garrett:2010hd}, for a recent review, and references therein.} One of the interesting candidates for DM is a keV sterile neutrino since some astrophysical data possibly support the existence~\\cite{Kusenko:2009up,Kusenko}. In extentions of the SM, the right-handed neutrinos can be added. The models with three right-handed neutrinos whose masses are below the EW scale have been proposed in~\\cite{deGouvea:2005er,Asaka:2005an,Asaka:2005pn}. In particular, the ref.~\\cite{Asaka:2005pn} has pointed out that the model with one keV sterile (right-handed) neutrino and two GeV range of ones can explain simultaneously the DM and baryon asymmetry of the Universe (BAU) in addition to the results of neutrino oscillation experiments. This framework is known as $\\nu$MSM \\cite{Asaka:2005an,Asaka:2005pn}. Furthermore, the keV sterile neutrino has been also considered in a gauge extension of the SM (left-right symmetric framework) \\cite{Bezrukov:2009th}. A realization of splitting pattern of right-handed neutrino mass spectrum including keV mass state in the context of flavor symmetry has been also discussed in \\cite{Shaposhnikov:2006nn,Lindner:2010wr,Merle:2011yv}. In \\cite{Shaposhnikov:2006nn}, it has been shown that the degeneracy of two singlet Majorana neutrinos and the lightness (keV scale) of the third one in $\\nu$MSM can be a consequence of a lepton number symmetry, which is broken in both the Yukawa and Majorana mass sectors. In \\cite{Lindner:2010wr}, a keV sterile neutrino as a candidate for DM is induced from softly breaking of $L_e-L_\\mu-L_\\tau$ flavor symmetry. In \\cite{Merle:2011yv}, a realization of the keV sterile neutrino arises from the Froggatt-Nielsen mechanism \\cite{Froggatt:1978nt}. Other interesting direction to realize the keV sterile neutrino DM and BAU has been discussed in the split seesaw mechanism \\cite{Kusenko:2010ik}. This mechanism can realize a splitting mass spectrum of the right-handed neutrinos in the context of extra-dimension and account for the smallness of active neutrino masses without fine-tuning. In the mechanism, the lightest right-handed neutrino with keV mass becomes DM candidate and the other two right-handed neutrinos make a source of BAU via leptogenesis \\cite{Fukugita:1986hr}.  Regarding with properties of active neutrinos, various neutrino oscillation experiments have shown that the neutrino mixing pattern is peculiar, namely, there are two large mixing angles and one small one~\\cite{Schwetz:2011qt}. A lot of works have been proposed to explain such a peculiar mixing pattern, e.g. $\\mu-\\tau$ symmetric texture~\\cite{Fukuyama:1997ky}, tri-bimaximal mixing~\\cite{Harrison:2002er}, golden ratios~\\cite{goldenratio}, and etc.. A proposed mixing pattern can be explained by imposing a non-Abelian discrete flavor symmetry with its breaking in a specific direction. For example, $\\mu-\\tau$ symmetry can be explained by the group $S_3 \\simeq D_3$~\\cite{D3} or $D_4$~\\cite{D4}. The tri-bimaximal mixing can be obtained from $A_4$~\\cite{A4} or $S_4$~\\cite{Ma:2005pd}. The golden ratios have been discussed with $A_5$~\\cite{A5} and $D_{10}$~\\cite{D10}.\\footnote{Our references for flavor models based on non-Abelian discrete symmetry are not complete. See e.g.~\\cite{Altarelli:2010gt} for more complete list.} Such the non-Abelian discrete symmetries might arise from the breaking of the gauge symmetries~\\cite{ConSym} or from the orbifold compactification in extra-dimensions~\\cite{OrbifoldSym}. Moreover, the vacuum alignment of the flavon field can be achieved without dealing with a relatively complicated scalar potential for $A_4$~\\cite{A4breaking} and $S_4$~\\cite{Adulpravitchai:2010na} symmetries in the extra-dimensional theory.  In this paper, we study models based on a flavor symmetry in the context of the split seesaw mechanism which can lead to a splitting mass spectrum of the right-handed neutrinos.\\footnote{The ref.~\\cite{Barry:2011wb} has studied the realization of eV scale sterile neutrinos within both seesaw mechanism and flavor symmetry. The work has shown that light sterile neutrinos can be accommodated in $A_4$ flavor models.} The paper is organized as following: In section 2, a brief review of the split seesaw mechanism is given. In section 3, we discuss the $A_4$ flavor models in the split seesaw mechanism. Then it will be shown how the desired flavon vacuum alignment can be obtained. We also give a comment on the leptogenesis. Section 4 is devoted to the conclusion. ",
        "conclusions": "We have studied models based on $A_4$ flavor symmetry in the context of the split seesaw mechanism. The split seesaw mechanism, which has been proposed an extra-dimensional theory, can provide a natural way to realize a splitting mass spectrum of right-handed neutrinos. It leads to one keV sterile neutrino as a DM candidate favored by recent cosmological observations and two heavy right-handed neutrinos being responsible for leptogenesis to explain the observed BAU.  First, we have pointed out that most of known $A_4$ flavor models with three right-handed neutrinos being $A_4$ triplet suffer from a {\\it degeneracy problem} for the bulk mass term, which disturbs the split mechanism for right-handed neutrino mass spectrum.  Next, we have constructed a new $A_4$ flavor model to work in the split seesaw mechanism. The model is an extention of the model proposed by Ma. The three right-handed neutrinos are assigned into $A_4$ (non-trivial) singlets, and the neutrino masses and mixing angles can be realized from both type I$+$II seesaw contributions in the model. The model predicts the $\\mu-\\tau$ symmetry in the neutrino mass matrix at the leading order, resulting in the vanishing $\\theta_{13}$ and maximal $\\theta_{23}$. The flavor symmetry $A_4$ is broken via the flavon vacuum alignment which can be obtained from the orbifold compactification. The BAU can also be realized in the mixed scenario of type I$+$II leptogenesis.  \\subsection*{Acknowledgement}  We are grateful to Takehiko Asaka, Peihong Gu, Claudia Hagedorn, Andreas Hohenegger, Martin Holthausen, Alexander Kartavtsev, Manfred Lindner, Kristian McDonald, and Werner Rodejohann for fruitful discussions. This work has been supported by the DFG Sonderforschungsbereich Transregio 27 Neutrinos and beyond Weakly interacting particles in Physics, Astrophysics and Cosmology.  \\appendix"
    },
    "1107/1107.3368_arXiv.txt": {
        "abstract": "The estimation of molecular abundances in interstellar clouds from spectroscopic observations requires radiative transfer calculations, which depend on basic molecular input data. This paper reviews recent developments in the fields of molecular data and radiative transfer. The first part is an overview of radiative transfer techniques, along with a \"road map\" showing which technique should be used in which situation. The second part is a review of measurements and calculations of molecular spectroscopic and collisional data, with a summary of recent collisional calculations and suggested modeling strategies if collision data are unavailable. The paper concludes with an overview of future developments and needs in the areas of radiative transfer and molecular data. ",
        "introduction": "\\label{s:intro}  At the heart of observational astrochemistry is the estimation of molecular abundances from telescope data. Extracting the information out of the data requires radiative transfer calculations which need molecular input data, and these two topics are discussed in this paper. In many cases, measuring the chemical composition of interstellar gas is coupled to the estimation of its physical conditions, in particular the kinetic temperature and the volume density. In addition, molecular abundances usually need interpretation with a chemical model, to estimate parameters such as the age of the cloud or its ionization rate by cosmic rays. These two important related topics are not covered here.  Before starting the review, it is useful to consider the limitations of radiative transfer calculations. Some limitations are of observational nature, such as resolution mismatch between observations at different wavelengths and the presence of unresolved substructure in the data. In addition, the lines in a spectrum may not all sample the same gas, for instance when some lines appear in emission and others in absorption. Furthermore, the observations may only cover part of the source, and only a few molecular energy levels. Finally the radiative transfer calculations may be limited by the availability of molecular data: either spectroscopic information or  collisional excitation probabilities. When interpreting observations, it is important to keep these limitations in mind and to choose the model such to minimize their impact.  The field of molecular data and radiative transfer is large, and due to space limitations, this paper cannot treat every aspect in depth. This review focuses on the analysis of spectral lines; see \\citet{pascucci2004} for a discussion of continuum radiative transfer and \\citet{shirley2011} for a recent update about dust opacities. In the area of line radiative transfer, see \\citet{peraiah2001} for an introduction and \\citet{pavlyuchenkov2007} for recent developments. For detailed accounts on laboratory spectroscopy, see \\citet{pickett1998} and \\citet{mueller2005} and references therein. The book by \\citet{flower2003} reviews the field of interstellar molecular collisions. Finally, the estimation of physical parameters of interstellar clouds is reviewed by \\citet{evans1999}. Note that throughout this review, the term `molecules' includes mono-atomic and ionic species. ",
        "conclusions": ""
    },
    "1107/1107.1471_arXiv.txt": {
        "abstract": "We report {\\it Fermi}/LAT observations and broad-band spectral modeling of the radio-loud active galaxy 4C\\,+55.17 (z=0.896), formally classified as a flat-spectrum radio quasar.  Using 19 months of all-sky survey {\\it Fermi}/LAT data, we detect a $\\gamma$-ray continuum extending up to an observed energy of $145$\\ GeV, and furthermore we find no evidence of $\\gamma$-ray variability in the source over its observed history. We illustrate the implications of these results in two different domains.  First, we investigate the origin of the steady $\\gamma$-ray emission, where we re-examine the common classification of 4C\\,+55.17 as a quasar-hosted blazar and consider instead its possible nature as a young radio source. We analyze and compare constraints on the source physical parameters in both blazar and young radio source scenarios by means of a detailed multiwavelength analysis and theoretical modeling of its broad-band spectrum. Secondly, we show that the $\\gamma$-ray spectrum may be formally extrapolated into the very-high energy (VHE; $\\geq 100$\\,GeV) range at a flux level detectable by the current generation of ground-based Cherenkov telescopes. This enables us to place constraints on models of extragalactic background light (EBL) within LAT energies and features the source as a promising candidate for VHE studies of the Universe at an unprecedented redshift of z=0.896. ",
        "introduction": "\\label{section-intro}  The radio-loud active galaxy 4C\\,+55.17 (0954+556), formally classified as a flat spectrum radio quasar (FSRQ), has a history of $\\gamma$-ray observations dating back to the EGRET era, as 3EG\\,J0952+5501 \\citep{har99,mat01} and EGR\\,J0957+5513 \\citep{cas08}. Due to a relatively poor localization of the EGRET source, however, the association of the $\\gamma$-ray emitter with 4C\\,+55.17 remained tentative at that time. After the successful launch of the {\\it Fermi} Gamma-Ray Space Telescope in June 2008, this association was on the other hand quickly confirmed by the Large Area Telescope \\citep[LAT;][]{atw09}, initially as 0FGL\\,J0957.6+5522 \\citep{bsl,lbas}, and most recently as 1FGL\\,J0957.7+5523 \\citep{1fgl}.  The quasar classification of 4C\\,+55.17 may be attributed to the presence of broad optical emission lines in its spectrum \\citep{wil95} and high optical/UV core luminosity \\citep[absolute $B$-band magnitude, $M_{\\rm B} < -23$;][]{ver06}.  Its redshift\\footnote{Assuming a $\\Lambda$CDM cosmology with $H_{\\rm 0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$, and $\\Omega_{\\rm \\Lambda}=0.73$, the luminosity distance $d_{\\rm L}=5785$\\,Mpc, and the conversion scale is $1$\\,mas\\,$=7.8$\\,pc.}, $z=0.896$, is based on the detection of Ly$\\alpha$ and CIV lines with the HST-FOS \\citep{wil95} and Mg II in the SDSS spectrum \\citep{sch07}.  The optical-UV properties of the source, together with its high $\\gamma$-ray luminosity of the order $L_{\\gamma} \\simeq 10^{47}$\\,erg\\,s$^{-1}$, have in turn led to the common classification of 4C\\,+55.17 as a blazar/FSRQ.  However, 4C\\,+55.17 also exhibits a number of morphological and spectral properties that have placed its exact blazar/FSRQ classification into question \\citep{mar02,ros05}.  FSRQs are uniquely characterized by the presence of a central compact radio core exhibiting a highly variable flat-spectrum continuum, high brightness temperatures ($T_{\\rm b}$), and, typically, superluminal motions on VLBI scales \\citep{urry95}.  Indeed, all of the aforementioned radio properties are shared by the luminous blazars detected in $\\gamma$-rays: these are exclusively observed to possess compact, highly polarized jets a few milli-arcseconds (mas) in angular size, and unresolved radio cores with brightness temperatures in the range $T_{\\rm b}=10^{10}-10^{14}$\\,K when observed at 5 GHz \\citep{tay07} and 15 GHz \\citep{kov09}.  In comparison, 4C\\,+55.17 demonstrates none of these characteristics. To date, the source shows no evidence of blazar flaring at any wavelength, nor any evidence of long-term variability, with the exception of a $\\sim 30\\%$ optical flux-density change noted over a period of 7 years between recent {\\it Swift}/UVOT measurements and archival SDSS data (see \\S\\,\\ref{section-other_data} and \\S\\,\\ref{section-model} for discussion). Furthermore, the VLBI radio morphology of the source is extended over $\\sim 400$\\,pc (projected). The peak surface brightness in a VLBA 15 GHz image taken from \\citet{ros05} is found in the northernmost component and is clearly resolved, with a corresponding brightness temperature $T_{\\rm b} < 2 \\times 10^{8}$ K \\citep[consistent with a measurement at 5 GHz;][]{tay07}, which is uncharacteristic of all the other known quasar-hosted $\\gamma$-ray blazars.  Based on the radio morphology of 4C\\,+55.17, \\cite{ros05} first suggested that the source may in fact belong to the family of young radio sources \\citep[for a review, see][]{odea98}, rather than blazars. Such sources are characterized by a very low radio variability (if any) and symmetric double radio structures resembling ``classical doubles'' on much smaller scales: linear sizes $\\lesssim 1$\\,kpc for compact symmetric objects (CSOs) and $\\sim 1-15$\\,kpc for medium symmetric objects \\citep[MSOs;][]{aug06}, to be compared with the typical linear sizes of ``regular'' Fanaroff-Riley type-II radio galaxies of $\\sim 100$\\,kpc. In many cases, CSO sources are found to exhibit a turnover in their radio spectra in the range of $0.5 - 10$\\,GHz, as the so-called \u201cGigahertz Peaked Spectrum (GPS) objects\u201d do \\citep{dev97}; similarly, MSO's often display turnover frequencies below $0.5$\\,GHz, typical of the Compact Steep Spectrum (CSS) class of sources \\citep{fan90}.  The overlap between CSO and GPS samples, as well as between samples of MSOs and CSS sources, is however not complete \\citep{sne00,aug06}.  In the case of 4C\\,+55.17, the VLBI morphology at $5$\\,GHz reveals two distinct emission regions, to the north and south \\citep[][see also Figure\\,\\ref{figure-radio}]{ros05}, covering a total angular extent of $53$\\,mas ($=413$\\,pc, projected). On the kpc scale, the source reaches $4\\arcsec.5$ ($\\sim 35$\\,kpc, projected), and it is resolved with the VLA in three components, the central one hosting the VLBI structure.  The northern component of the pc-scale emission features a compact region with a relatively flat spectrum \\citep[$\\alpha=0.4; F_{\\nu}\\propto\\nu^{-\\alpha}$;][]{ros05}, which can be attributed to a core or a hotspot region, while the southern component features a more diffuse and slightly steeper-spectrum ($\\alpha=0.49$) region.  \\citet{ros05} have pointed out that these two components resemble more compact hotspots and lobes, suggesting a CSO/MSO classification for this object. The kpc-scale emission might thus be interpreted as a remnant of previous jet activity, as this is a common feature among sources that show evidence of intermittent behavior \\citep[e.g.,][]{baum90,luo07,ori08}.  Under the CSO/MSO framework, \\citeauthor{ros05} found no core candidate between the VLBA-scale lobes at a level $\\gtrsim 2$\\,mJy/beam in a $15$\\,GHz map.  An 11-month comparison of the $\\gamma$-ray variability and spectral properties of 4C\\,+55.17 against the other LAT FSRQs highlights the unusual nature of the source \\citep{1fgl,1lac}.  Among all of the sources originally detected in the 3-month LAT Bright AGN Sample \\citep[LBAS;][]{lbas} that were classified as FSRQs, 4C\\,+55.17 is characterized by the lowest variability index \\citep{1fgl}.  In addition, the unusually hard $\\gamma$-ray continuum (that is, with a low photon index $\\Gamma$) is found to be one of the hardest among FSRQs in the 1st LAT AGN Catalog \\citep[1LAC;][]{1lac}.  In fact, of those sources included in the 1LAC (FSRQ or otherwise) with $>1$\\,GeV flux greater than or equal to that of 4C\\,+55.17, only five -- all of which are BL Lac objects (PKS~2155--304, Mkn 421, 3C 66A, PG~1553+113, and PKS~0447-439) -- appear with a harder $\\gamma$-ray spectrum.  In this work, we re-examine the high-energy $\\gamma$-ray ($>100$\\,MeV) properties of 4C\\,+55.17 using 19 months of LAT all-sky survey data and discuss the implications of these results in two domains.  First, we reconsider the underlying physical processes responsible for the $\\gamma$-ray emission through detailed broadband modeling of the source in the context of two scenarios: ``young radio source'' and ``blazar.''  In addition, we demonstrate that the unusual properties of the source make it an ideal candidate for studying the high-redshift universe at very-high energies (VHE), in particular for placing constraints on the level of extragalactic background light (EBL).  The paper is organized as follows.  Section\\,\\ref{section-observations} details the analysis of 19 months of LAT data and discusses the supporting multiwavelength observations.  In particular, section\\,\\ref{section-lat} focuses on the LAT data reduction, presenting new spatial (localization), spectral, and variability analysis, including a detailed analysis of the 145 GeV photon detection associated with the source (see also Appendix\\ \\ref{section-photon}). Section \\ref{section-other_data} discusses the multiwavelength observations, including analysis of archival radio and {\\it Swift} X-ray and optical data, as well as a new hard X-ray detection with the {\\it Swift} Burst Alert Telescope (BAT).  Spectral properties and classification of 4C\\,+55.17 are discussed in section\\,\\ref{section-discussion}. We follow with a detailed analysis of the high energy spectrum of 4C\\,+55.17, where we place constraints on models of EBL and discuss the implications of the 145 GeV photon detection to future VHE observations of the source (\\ \\S\\,\\ref{section-EBL}).  Our conclusions are presented in section\\,\\ref{section-conclusions}. ",
        "conclusions": "The investigation of the multiwavelength properties of 4C\\,+55.17, including its unusually hard $\\gamma$-ray spectrum, lack of distinct variability, and CSO-like radio morphology, has highlighted the exceptional nature of this $\\gamma$-ray source.  For the first time, we have modeled the radio to $\\gamma$-ray emission of 4C\\,+55.17 as a young radio source using a dynamic model that is consistent with the full extent of its observed properties. Furthermore, we find that the prospect of a VHE observation of 4C\\,+55.17, whose $\\gamma$-ray spectrum already extends up to the observed energy of 145\\ GeV, is within reach of the current generation of Cherenkov telescopes.  A detection by such an instrument would not only add to the present understanding of the source itself, but would also serve to place a significant upper limit to the level of EBL through combined {\\it Fermi} and VHE data.  Furthermore, we anticipate that through continued monitoring of 4C\\,+55.17 at high energies with the {\\it Fermi} LAT, as well as in the radio through X-rays, the precise classification of 4C\\,+55.17 will become increasingly more apparent.  If, for example, the source continues to remain non-variable in $\\gamma$-rays in the years to come, its average flux versus variability may eventually lie outside the distribution of {\\it Fermi} $\\gamma$-ray emitting blazars altogether, which would make any standard blazar emission scenario difficult to reconcile.  On the other hand, if rapid variability is found in this source, that would seem to rule out a pure CSO interpretation. Thus we expect that 4C\\,+55.17 will be an important target for future observations across all wavelengths."
    },
    "1107/1107.0899_arXiv.txt": {
        "abstract": "Although the rapid neutron-capture process, or r-process, is fundamentally important for explaining the origin of approximately half of the stable nuclei with $A > 60$, the astrophysical site of this process has not been identified yet. Here we study r-process nucleosynthesis in material that is dynamically ejected by tidal and pressure forces during the merging of binary neutron stars (NSs) and within milliseconds afterwards. For the first time we make use of relativistic hydrodynamical simulations of such events, defining consistently the conditions that determine the nucleosynthesis, i.e., neutron enrichment, entropy, early density evolution and thus expansion timescale, and ejecta mass. We find that $10^{-3}$--$10^{-2}\\,M_\\odot$ are ejected, which is enough for mergers to be the main source of heavy ($A \\gtrsim 140$) galactic r-nuclei for merger rates of some $10^{-5}$yr$^{-1}$. While asymmetric mergers eject 2--3 times more mass than symmetric ones, the exact amount depends weakly on whether the NSs have radii of $\\sim$15\\,km for a ``stiff'' nuclear equation of state (EOS) or $\\sim$12\\,km for a ``soft'' EOS. R-process nucleosynthesis during the decompression becomes largely insensitive to the detailed conditions because of efficient fission recycling, producing a composition that closely follows the solar r-abundance distribution for nuclei with mass numbers $A > 140$. Estimating the light curve powered by the radioactive decay heating of r-process nuclei with an approximative model, we expect high emission in the B-V-R bands for 1--2 days with potentially observable longer duration in the case of asymmetric mergers because of the larger ejecta mass. ",
        "introduction": "The r-process, or rapid neutron-capture process, of stellar nucleosynthesis is invoked to explain the production of the stable (and some long-lived radioactive) neutron-rich nuclides heavier than iron, which are observed in stars of different metallicities as well as in the solar system \\citep[for a review, see][]{ar07}. Despite a growing wealth of observational data \\citep[e.g.,][]{sneden08} and although increasingly better r-process models with new astrophysical or nuclear physics ingredients have been developed over decades, the stellar production site(s) of r-process material has (have) not been identified yet. All proposed scenarios face serious problems. Supernovae (SNe), for example, appear attractive because of their potential to explain observational features of the galactic chemical enrichment history \\citep[e.g.,][]{argast04}. Their nucleosynthesis, however, exhibits extreme sensitivity to the detailed conditions in the ejecta, whose viability for strong r-processing could not be verified by sophisticated hydrodynamical models \\citep[e.g.,][]{hoffman08,janka08,roberts10,hudepohl10,fisch10,wanajo11}.   Early in the development of the theory of nucleosynthesis, an alternative origin of r-process nuclei was proposed \\citep{tsuruta65}. It relies on the fact that at high densities (typically $\\rho > 10^{10}$ gcm$^{-3}$) matter tends to be composed of nuclei lying on the neutron-rich side of the valley of nuclear stability as a result of endothermic free-electron captures. The astrophysical plausibility of this production mechanism as the source of the observed r-nuclides has long been questioned. It remained largely unexplored until \\citet{lattimer77} and \\citet{meyer89} studied the decompression of cold, neutronized matter ejected by tidal effects of a black hole (BH) on a neutron star (NS) companion. Recently, special attention has been paid to NS mergers because hydrodynamic simulations of NS-NS and NS-BH mergers showed that a non-negligible amount of matter may be ejected \\citep[e.g.,][]{janka99,ros04,oech07}.  All previous investigations of the ejecta from coalescing NSs \\citep{frei99,go05,ar07,metzger10,roberts11} were parametrized in one way or another, which makes their results and conclusions subject to open questions. In \\citet{go05} and \\citet{ar07} the thermodynamic profiles were constructed by a simple decompression model \\citep[see][]{go11}, but the neutron enrichment (or equivalently the electron fraction $Y_e$) was consistently taken from $\\beta$-equilibrium assumed to have been achieved at the initial density prior to the decompression. It was found that the final composition of the material ejected from the inner crust depends on the initial density, at least for the outer parts of the inner crust at $\\rho_{\\mathrm{drip}} \\le \\rho \\le 10^{12}$g/cm$^3$ (where $\\rho_{\\mathrm{drip}}\\simeq$4.2$\\times$10$^{11}$g/cm$^3$ is the neutron-drip density). For the deeper inner-crust layers ($\\rho > 10^{12}~{\\rm g/cm^{3}}$), large neutron-to-seed ratios  drive the nuclear flow into the very heavy-mass region, leading to multiple fission recycling. As a consequence, the resulting abundance distribution becomes independent of the initial conditions, especially of the initial density. It was found to be in close agreement with the solar distribution for $A>140$ nuclei \\citep{go05,ar07}.  Different approaches were taken to nucleosynthesis calculations for merger ejecta by \\citet{frei99,metzger10} and \\citet{roberts11}. In their calculations, while the density evolution of the mass elements was adopted from hydrodynamical simulations, both the initial neutron enrichment and the temperature history were considered as free parameters. In particular, $Y_e$ was chosen in order to obtain, after decompression, an r-abundance distribution as close as possible to the solar distribution. This led to values of $Y_e=0.1$ \\citep{frei99} or 0.2 \\citep{roberts11}, corresponding to relatively near-surface layers of the inner crust and to nuclear flows that are not subject to multiple fission cycles. In our simulations, most of the ejecta mass originates from the deep layers of the inner crust so that any contribution from near-surface layers remains minor. In addition, previous studies assumed that the nucleosynthesis is independent of the initial temperature of the ejected material, while we find that the initial temperature may not only affect the initial composition, but potentially also the nucleosynthesis (see Sect.~\\ref{sect_in}) so that special attention should be paid to the detailed temperature history of the ejected material prior to its free expansion.  The work presented here is based on recent three-dimensional relativistic simulations of NS-NS mergers to determine the nucleosynthesis-relevant conditions of the ejected matter. The hydrodynamical model is described in Sect.~\\ref{sect_hydro}. Section \\ref{sect_in} presents the nucleosynthesis results for two binaries, a symmetric NS-NS system and an asymmetric one. In contrast to previous studies, detailed information about the density, $Y_e$, and entropy evolution of the ejecta is extracted from the hydrodynamical simulations and included in the network calculations. The expected electromagnetic emission that is powered by radioactive decays following the heavy-element nucleosynthesis is estimated by a simple, approximative light-curve model in Sect.~\\ref{sect_lc}. Conclusions are drawn in Sect.~\\ref{sect_conc}.  \\begin{figure*} \\begin{center} \\includegraphics[width=2.\\columnwidth]{fig_histo.pdf} \\end{center} \\vskip -0.5cm \\caption{Histograms of fractional mass distribution of the ejecta for the 1.35--1.35\\Msun\\ NS merger (upper row) and the 1.2--1.5\\Msun\\ binary (lower row) as functions of density $\\rho$ (relative to the saturation density $\\rho_S\\simeq 2.6\\times$10$^{14}{\\rm g/cm^3}$; left) and of electron fraction $Y_e$ (middle) that the ejected matter had at its initial NS location prior to merging. The right panels show the fractional mass distributions as functions of the final entropy $S$ per nucleon when the matter starts its free expansion. In the insert on the left panels the dots mark positions of mass elements that get ejected later. The locations are given in projection on the orbital plane at the time when the stellar collision begins. } \\label{fig_histo} \\end{figure*} ",
        "conclusions": "\\label{sect_conc}  Using relativistic NS merger models to determine the nucleosynthesis-relevant conditions self-consistently, we confirm that decompressed NS matter ejected dynamically during the stellar collision and shortly afterwards is an extremely promising site for robust, strong r-processing. Matter from the inner crust of the coalescing NSs, which dominates the ejecta by far, produces an r-abundance distribution very similar to the solar one for nuclei with $A>140$. Nuclei with $A<140$ with solar distribution could originate from the outer crust \\citep{go11}, but too little of such matter gets ejected to explain the solar proportion of light-to-heavy r-process material. However, significant amounts of $A < 140$ nuclei might be produced in the outflow of a BH-torus system formed after the NS merger (Wanajo \\& Janka 2011).  The underlying nuclear mechanisms differ significantly from those at action in SN scenarios. In particular fission plays a major role in recycling heavy material. The similarity between predicted and solar abundance patterns as well as the robustness of the prediction against variations of input parameters, which we have shown in \\cite{go05}, have demonstrated here in comparison of symmetric and asymmetric NS-NS mergers, and will further elaborate on in a forthcoming paper, make this site one of the most promising, deserving further exploration concerning hydrodynamics, galactic chemical evolution, nucleosynthesis, nuclear physics, and astronomical consequences.  Fully relativistic simulations including neutrino transport and magnetic fields, with good resolution of the inner (and outer) crust layers of the merging NSs assured by adaptive refinement, are needed to corroborate the ejecta conditions found in our work. With our yield of $\\sim$(3--5)$\\times 10^{-5}\\,M_\\odot$ per event of $^{151,153}$Eu, a nearly-pure r-process element, the origin of all galactic Eu from NS-NS mergers would require an event rate of $\\sim$(2--3)$\\times 10^{-5}$yr$^{-1}$, fully compatible with present best estimates \\citep[see, e.g.,][]{bel08}. However, also mass loss during the evolution of the merger remnant from a hypermassive NS to a massively accreting BH-torus system as well as NS-BH mergers with higher ejecta masses might add to the production (e.g., Wanajo \\& Janka 2011, Caballero et al.\\ 2011) and deserve more detailed investigations. From the nuclear point of view, this site also implies new challenges because it involves the formation of neutron-drip nuclei, for which $\\beta$-decay, neutron capture, and fission rates need to be determined. Astronomically, the discovery of electromagnetic radiation from kilo-nova events \\citep{roberts11,metzger10} could mean the first in-situ observation of freshly produced r-process material."
    },
    "1107/1107.3722.txt": {
        "abstract": "{We study the production of cosmogenic neutrinos and photons during the extragalactic propagation of ultra-high-energy cosmic rays (UHECRs). For a wide range of models in cosmological evolution of source luminosity, composition and maximum energy we calculate the expected flux of cosmogenic secondaries by normalizing our cosmic ray output to experimental spectra {and comparing the diffuse flux of GeV-TeV gamma-rays to the experimental one measured by the Fermi satellite}. Most of these models yield significant neutrino fluxes for current experiments like IceCube or Pierre Auger. Furthermore, we discuss the possibilities of signing the presence of UHE proton sources either within or outside the cosmic ray horizon using neutrinos or photons observations even if the cosmic ray composition becomes heavier at the highest energies. We discuss the possible constraints that could be brought on the UHECR origin from the different messengers and energy ranges. % ",
        "introduction": "\\label{Introduction}  The origin of ultra-high-energy cosmic rays is one of the main topics in high-energy astrophysics. Their origin still remains unknown after decades of experimental efforts. Recent experiments such as the High Resolution fly's eye (HiRes, Abbasi et al., 2004), the Akeno Giant Air Shower Array (AGASA, Nagano et al., 1992) and primarily the Pierre Auger Observatory  (Abraham et al., 2004) have increased the available cosmic ray statistics above $10^{18}$ eV, allowing the first solid studies from the point of view of the spectrum, the composition and the arrival directions of ultra-high-energy cosmic rays (UHECRs). A sharp decrease of the UHECR flux above 3-4$\\times 10^{19}$ eV possibly related to the theoretically expected GZK cut-off (Greisen, 1966; Zatsepin \\& Kuzmin, 1966) seems to be firmly established, whereas the interpretation of composition or arrival direction data as well as the consistency of the results reported by the different experiments are matters of intense debates.  A multimessenger approach of the UHECR question is a promising way for solving the mystery of their origin. Indeed, shortly after the prediction of a cut-off at the highest energies caused by the interactions of UHECR protons and nuclei with the cosmic microwave background (CMB) photons, the unavoidable associated production of UHE neutrinos through the decay of charged pions was pointed out by Berezinsky and Zatsepin (1969). Because these {\\it cosmogenic neutrinos} can travel from their production site without undergoing interactions or deflection, they were, soon after this pioneering work, considered as potentially interesting probes of the highest energy phenomena in the universe. Their expected flux on earth was intensively calculated in the past decades for different astrophysical scenarios on the cosmological evolution of the sources, the composition or the maximum energy at the sources (see for instance Stecker ,1979; Hill \\& Schramm,1985; Engel et al., 2001; Kalashev et al., 2002; Secker \\& Stanev, 2005; Hooper et al., 2005; Ave et al., 2005; Stanev et al., 2006; Allard et al., 2006; Anchordoqui et al., 2007; Takami et al., 2009; Berezinsky, 2009; Ahlers et al., 2009; Kotera et al., 2010).  Cosmogenic gamma-rays are also produced during the propagation of UHECRs. Unlike neutrinos, these very high-energy gamma-rays interact rapidly and produce electromagnetic cascades. As a result the universe is opaque to gamma-rays from a few hundreds of GeV to a few $10^{18}$ eV. Above $10^{19}$ eV the universe becomes more and more transparent to photons and very high-energy gamma-rays can propagate a few tens of megaparsecs without losing a great amount of energy. As a result, these very high-energy cosmogenic gamma-rays were discussed in the literature either as signatures of the so-called Top-Down models (see for instance  Protheroe \\& Johnson, 1996; Lee, 1998; Sigl et al., 1999; Semikoz \\& Sigl., 2004) or as probes of UHECR acceleration in the local universe (e.g, Yoshida \\& Teshima, 1993; Protheroe and Johnson, 1996; Lee, 1998; Semikoz \\& Sigl., 2004; Gelmini et al., 2007a and 2007b, Taylor \\& Aharonian, 2009; Taylor et al., 2009; Kuempel et al., 2009; Hooper et al., 2011; Ahlers \\& Salvado, 2011).  Because electromagnetic cascades, piling up below 100 GeV, are produced during UHECR propagation and are likely to be associated with the production of cosmogenic neutrinos, it has soon been realized that measurements of the diffuse gamma-ray background could allow one to put constraints on the cosmological evolution of the UHECR luminosity (Strong et al., 1973, 1974) and on the maximum allowable cosmogenic neutrino fluxes (Berezinsky \\& Smirnov, 1975). Modern versions of this calculation were attempted using the EGRET measurements (Sreekumar et al., 1998; Strong et al., 2004)  in Kalashev et al., (2002), Semikoz \\& Sigl (2004), Kalashev et al. (2007). More recently, the Fermi satellite measurements  (Abdo et al., 2010) reported a gamma-ray background between 100 MeV and 100 GeV lower than previously estimated using EGRET data. The additional constraints allowed by this new measurement were discussed by Berezinsky et al. (2010), Ahlers et al. (2010) and more recently by Wang et al. (2011). These studies agree that the constraints implied by this new estimate of the gamma-ray background are much more stringent than those using EGRET data, but their conclusions differ on the impact of these new constraints on the observability of UHE neutrino fluxes (see below). In this paper we study the production of cosmogenic photons and neutrinos assuming different astrophysical scenarios for the cosmological evolution of the source luminosity, the composition and different scenarios for the transition from galactic to extragalactic cosmic rays in the same spirit as in Kotera et al. (2010). We discuss the constraints brought by the Fermi measurements of the diffuse extragalactic background on the expectations for cosmogenic neutrinos or UHE photons detection. Furthermore, we discuss the possible neutrino or photon signatures from UHECR proton accelerators either in the local universe or at cosmological distances (diffuse or point sources) even if the UHECR composition becomes heavier, which is one of the plausible interpretations of the recent measurement of the longitudinal development of air showers (hereafter $X_{max}$ measurements) of the Pierre Auger Observatory (Abraham et al. 2010a). In the next section, we briefly introduce our code that simulates the propagation of UHECR protons and nuclei and our Monte Carlo tool for the development of intergalactic electromagnetic cascades. In section 3 we then calculate cosmogenic photons and neutrino fluxes by normalizing our simulations with the experimental spectrum measured by the Pierre Auger Observatory (Abraham et al., 2010b).  We finally discuss these results and the possible constraints on the UHECR origin that can potentially be achieved by observing either point sources or diffuse contributions of cosmogenic secondaries.   %\\begin{figure*}[ht] %\\centering{\\includegraphics[width=\\textwidth]{8653fig1.pdf}} %\\caption{Structure of a purely turbulent magnetic field with variance $\\sqrt{\\langle \\|\\mathbf{\\vec{B}^2}\\| \\rangle} = 10$~nG and maximum turbulence scale $\\lambda_{\\mathrm{max}} = 1$~Mpc. From left to right: distribution of the magnetic field squared norm, $\\|\\mathbf{\\vec{B}^2}\\|$, distribution of component $B_x$, and magnetic field two-point autocorrelation as a function of point separation.} %\\label{fig:fieldStructure} %\\end{figure*} ",
        "conclusions": "We have considered the production of cosmogenic secondary photons and neutrinos during the extragalactic propagation of UHE protons and nuclei. We discussed the constraints obtained from {the Fermi} observations of the diffuse gamma-ray flux. We found that, because the possibility of a high contribution of UHECR to the diffuse gamma-ray flux is currently not ruled out, significant UHECR neutrino fluxes observable by present observatories such as IceCube or the Pierre Auger Observatory can still be expected and would be of prime interest for the field. The UHE cosmogenic photons could also be expected for optimistic astrophysical assumptions. These fluxes are not constrained by the experimental diffuse gamma-ray flux because they show little dependance on the evolution of the source luminosity evolution unlike UHE neutrinos. We also discussed the influence of the composition evolution at the highest energies. We pointed out that scenarios for which the maximum energy at the source is limited (well below $10^{20}$ eV per nucleon) are expected to have their UHECR neutrino flux strongly suppressed above $10^{17}$ eV. However, we found that the likely variation of the maximum from source to source could yield an observable neutrino flux, if rare powerful accelerators (e.g, source able to accelerate particles above $10^{20}$ eV per nucleon) show a strong cosmological evolution of their luminosity and contribute at the level of $\\sim10\\%$ to the cosmic ray flux around $10^{19}$~eV.  We also considered the possibility of observing cosmogenic secondaries from individual sources to constrain the origin of highest energy particles. Distant and very powerful sources were found to be prime candidates, and neutrino flux along with synchrotron images (Gabici and Aharonian, 2005) would be unambiguous signatures of the acceleration of particles above the threshold of pion production. These observations would not strongly constrain the source composition, though. But observations of sources in the local universe would require, with the current experimental capabilities, too high contributions of these individual sources to the UHE cosmic ray spectrum. This can be alleviated, however, with improved sensitivities (it would require point source sensitivities below $\\rm10^{-9}\\, GeV\\,cm^{-2}\\,s^{-1}$ above $10^{17}$ eV for neutrino observatories, for instance), which would at the same time decrease the luminosity requirements for distant sources.  GeV-TeV photons from electromagnetic cascade, taking advantage of the upcoming very sensitive CTA observatory, could be detected even in the local universe by assuming a straight line propagation of the electromagnetic cascades. However, their detectability strongly depends on the strength of the extragalactic magnetic field and sources in the local universe are likely to become very hard to detect if the latter exceeds $\\sim10^{-12}$ G in the local universe. For a positive detection strong constraints on the extragalactic magnetic field could of course be obtained as well as a precise location of the source -- a critical milestone in high-energy astrophysics. However, these cascades are not a signature of the acceleration of cosmic rays above the threshold of pion production or of the source composition. Consequently, counterparts (primarily from cosmic ray observatories) would be needed to fully understand the signal.  We conclude that the observation of cosmogenic secondaries either by their diffuse flux or individual sources can be expected from current or next-generation instruments. Importantly, they would represent invaluable additions to observations of the UHECR sky made by the Pierre Auger Observatory, the Telescope Array and the future JEM-EUSO in their quest for solving the long-standing UHECR problem and identifying its astrophysical source."
    },
    "1107/1107.2414.txt": {
        "abstract": "The secular approximation for the evolution of hierarchical triple configurations has proven to be very useful in many astrophysical contexts, from planetary to triple-star systems.  In this approximation the orbits may change shape and orientation, on time scales longer than the orbital time scales,  but the semimajor axes are constant.  For example, for highly inclined triple systems, the Kozai-Lidov mechanism can produce large-amplitude oscillations of the eccentricities and inclinations.  Here we revisit the secular dynamics of hierarchical triple systems.  We derive the secular evolution equations to octupole order in Hamiltonian perturbation theory.  Our derivation corrects an error in some previous treatments of the problem that implicitly assumed a conservation of the  z-component of the angular momentum of the inner orbit (i.e., parallel to the total angular momentum of the system). Already to quadrupole order, our results show new behaviors including the possibility for a system to  oscillate from prograde to retrograde orbits. At the octupole order, for an eccentric outer orbit, the inner orbit can reach extremely high eccentricities and undergo chaotic flips in its orientation.  We discuss applications to a variety of astrophysical systems, from stellar triples to merging compact binaries and planetary systems.  Our results agree with those of previous studies done to quadrupole order only in the limit in which one of the inner two bodies is a massless test particle and the outer orbit is circular; our results agree with previous studies at octupole order for the eccentricity evolution, but not for the inclination evolution. %  our results agree with %  previous studies at octupole order in the eccentricity evolution, %  but the inclination evolution we present has not been considered by %  previous studies. ",
        "introduction": "Triple star systems are believed to be very common \\citep[e.g.,][]{T97,Eggleton+07}. From dynamical stability arguments these must be hierarchical triples, in which the (inner) binary is orbited by a third body on a much wider orbit.  Probably more than 50\\% of bright stars are at least double \\citep{T97,Eggleton+07}. Given the selection effects against finding faint and distant companions we can be reasonably confident that the proportion is actually substantially greater.  \\citet{T97} showed that $40\\%$ of binary stars with period $<10$~d in which the primary is a dwarf ($0.5 -1.5\\,M_{\\odot}$) have at least one additional companion. He found that the fraction of triples and higher multiples among binaries with period ($10-100\\,$d) is $\\sim10\\%$. Moreover, \\citet{Pri+06} have surveyed a sample of contact binaries, and noted that among 151 contact binaries brighter than 10 mag., 42$\\pm5\\%$ are at least triple.  Many close stellar binaries with two compact objects are likely produced through triple evolution.  Secular effects (i.e., coherent interactions on timescales long compared to the orbital period), and specifically Kozai-Lidov cycling \\citep[][see below]{Kozai,Lidov}, have been proposed as an important element in the evolution of triple stars \\citep[e.g.][]{Har69,Mazeh+79,Sod82,1998KEM,Dan,PF09,Tho10,ST12}.  In addition, Kozai-Lidov cycling has been suggested to play an important role in both the growth of black holes at the centers of dense star clusters and the formation of short-period binary black holes \\citep{Wen,MH02,Bla+02}.  Recently, \\citet{Iva+10} showed that the most important formation mechanism for black hole XRBs in globular clusters may be triple-induced mass transfer in a black hole-white dwarf binary.  Secular perturbations in triple systems also play an important role in planetary system dynamics.  \\citet{Kozai} studied the effects of Jupiter's gravitational perturbation on an inclined asteroid in our own solar system. In the assumed hierarchical configuration, treating the asteroid as a test particle, \\citet{Kozai} found that its inclination and eccentricity fluctuate on timescales much larger than its orbital period. Jupiter, assumed to be in a circular orbit, carries most of the angular momentum of the system. Due to Jupiter's circular orbit and the negligible mass of the asteroid, the system's potential is axisymmetric and thus the component of the inner orbit's angular momentum along the total angular momentum is conserved during the evolution.  \\citet{Kozia79} also showed the importance of secular interactions for the dynamics of comets \\citep[see also][]{Quinn+90,Bailey+92,Thomas+96}.  The evolution of the orbits of binary minor planets is dominated by the secular gravitational perturbation from the sun \\citep{PN09}; properly accounting for the resulting secular effects---including Kozai cycling---accurately reproduces the binary minor planet orbital distribution seen today \\citep{Naoz10obs,Grundy+11}. In addition \\citet{Kin+91}, \\citet{vas99}, \\cite{Carr+02}, \\citet{Nes+03}, \\citet{cuk+04} and \\citet{KN07} suggested that secular interactions may explain the significant inclinations of gas giant satellites and Jovian irregular satellites.  Similar analyses have been applied to the orbits of extrasolar planets \\citep[e.g.,][]{Inn+97,Wu+03,Dan,Wu+07,Naoz10,Veras+10,Cor+11}. \\citet{Naoz10} considered the secular evolution of a triple system consisting of an inner binary containing a star and a Jupiter-like planet at several AU, orbited by a distant Jupiter-like planet or brown-dwarf companion. Perturbations from the outer body can drive Kozai-like cycles in the inner binary, which, when planet-star tidal effects are incorporated, can lead to the capture of the inner planet onto a close, highly-inclined or even retrograde orbit, similar to the orbits of the observed retrograde ``hot Jupiters.''  Many other studies of exoplanet dynamics have considered similar systems, but with a very distant stellar binary companion acting as perturber. In such systems, the outer star completely dominates the orbital angular momentum, and the problem reduces to test-particle evolution \\citep[see ][]{LN,Boaz2,Naoz+12}. If the lowest level of approximation is applicable (e.g., the outer perturber is on a circular orbit), the $z$-component of the inner orbit's angular momentum is conserved \\citep[e.g.,][]{LZ74}.   In early studies of high-inclination secular perturbations \\citep{Kozai,Lidov}, the outer orbit was circular and again dominated the orbital angular momentum of the system. In this situation, the component of the inner orbit's angular momentum along the z-axis is conserved.  In many later studies the assumption that the $z$-component of the inner orbit's angular momentum is constant was built into the equations \\citep[e.g.][]{1998EKH,Mik+98,Zdz+07}. %,Takeda}. In fact these studies are only valid in the limit of a test particle forced by a perturber on a circular orbit.  To leading order in the ratio of semimajor axes, the double averaged potential of the outer orbit is axisymmetric (even for an eccentric outer perturber), thus if taken to the test particle limit, this results in a conservation of the $z$-component of the inner orbit's angular momentum.  We refer to this limit as the ``standard'' treatment of Kozai oscillations, i.e. quadrupole-level approximation in the test particle limit (test particle quadrupole, hereafter TPQ).   In this paper we show that a common mistake in the Hamiltoniano treatment of these secular systems can lead to the erroneous conclusion that the $z$-component of the inner orbit's angular momentum is constant outside the TPQ limit; in fact, the $z$-component of the inner orbit's angular momentum is only conserved by the evolution in the test-particle limit and to quadrupole order. To demonstrate the error we focus on the quadrupole (non-test-particle) approximation in the main body of the paper, but we include the full octupole--order equations of motion in an appendix.  {In what follows we show the applications of these two effects (i.e., correcting the error and including the full octupole--order equations of motion) by considering different astrophysical systems. Note that the applications illustrated in the text are inspired by real systems; however, we caution that we consider here only Newtonian point mass dynamics, while in reality other effects such as tides and general relativity can greatly effect the evolution. For example, general relativity may alter the evolution of the system, which can give rise to a resonant behavior of the inner orbit\u00d5s eccentricity  \\citep[e.g.,][]{Ford00,Naoz12b}. Furthermore, tidal forces can suppress the eccentricity growth of the inner orbit, and thus significantly modify the evolutionary track of the system  \\citep[e.g.,][]{Mazeh+79,Sod84,1998KEM}. In particular tides, in some cases, can considerably suppress the chaotic behavior that arises in the presence of the the octupole\u00d0level of approximation  \\citep[e.g.,][]{Naoz10,Naoz+12}. Therefore, while the examples presented in this paper are inspired by real astronomical systems, the true evolutionary behavior will be modified from what we show once the eccentricity becomes too high. }  This paper is organized as follows. We first present the general framework (\\S \\ref{Sec:Ham}); we then derive the complete formalism for the quadrupole-level approximation and the equations of motion (\\S \\ref{sec:4eq}), we also develop the octupole-level approximation equations of motion in \\S \\ref{sec:H8}. We discuss a few of the most important implications of the correct formalism in \\S \\ref{sec:implications}. We also compare our results with those of previous studies (\\S \\ref{sec:implications}) and offer some conclusions in \\S \\ref{Sec:con}. ",
        "conclusions": "\\label{Sec:con}  We have shown that the ``standard'' TPQ Kozai formalism \\citep{Kozai,Lidov} has been applied in inappropriate situations.  A common error in the implementation of the relevant Hamiltonian mechanics (premature elimination of the nodes) leads to the (incorrect) conclusion that the conservation of the $z$-component of each orbit's angular momentum from the TPQ dynamics generalizes beyond the TPQ approximation.  Correcting the formalism we find that the $z$-components of  \\emph{both} the inner and outer orbits' angular momenta in general change with time at both the quadrupole and octupole level.  The conservation of the inner orbit's $z$-component of the angular momentum (the famous $\\sqrt{1-e_1^2}\\cos i={\\rm constant}$) only holds in the quadrupole-level \\emph{test particle} approximation.  We have explained in details the source of the error in previous derivations (Appendix \\ref{prob}).  We have re-derived the secular evolution equations for triple systems using Hamiltonian perturbation theory to the octupole-level of approximation (Section \\ref{Sec:Ham} and Appendix \\ref{Sec:form}, \\ref{sec:H8} and Appendix \\ref{sec:8}). We have also shown that one can use the simplified Hamiltonian found in the literature \\citep[e.g.,][]{Ford00} as long as the equations of motion for the inclinations are calculated from the total angular momentum.  The correction shown here has important implications to the evolution of triple systems. We discussed a few interesting implications in Section \\ref{sec:implications}. We showed that already at the quadrupole-level approximation the explicit assumption that the vertical angular momentum is constant can lead to erroneous results, see for example Figure \\ref{fig:quad}. In this Figure we showed that far from the test particle limit in the quadrupole-level one can already find a significant difference in the evolutionary behavior. The correct results agree with the test particle limit only when  $G_1/G_2< 10^{-4}$ (see Figure \\ref{fig:H1}).  We show in Appendix \\ref{sec:maxmin} that at the quadrupole level of approximation, the inner eccentricity and the mutual inclination have a well defined maximum and minimum irrespective of the mass of the inner bodies. In the test particle limit these values converge to the well-known critical inclinations ($39.2^\\circ \\leq i_0 \\leq 140.8^\\circ $) for large oscillatory amplitudes.  {The most notable outcome of the results presented here happens in the octupole-level of approximation (which we call the EKL formalism), when the inner orbit flips from prograde to retrograde with respect to the total angular momentum.  Just before the flip the inner orbit has an excursion of extremely high eccentricity. In the presence of tidal forces (not included in this study) the outcome of a system can be  different than the one assumed while using the TPQ formalism. \\citet{KM99}, \\citet{Ford00}, \\citet{Bla+02}, \\citet{Lee03} and \\citet{Las+10}  present the correct octupole equations of motion.  Had these authors integrated their equations for systems such as those presented in this paper, they could already have discovered the possibility of flipping the inner orbit. }  % %had the correct equations of motion, and could, in %principle, have observed this phenomena.  However, it seems that the %assumption of a constant vertical angular momentum was built into the %community understanding of Kozai mechanism, causing the flipping %effect of the EKL mechanism to be overlooked. % % In \\citet{Naoz10} we suggested that this effect may play an important % role in the formation mechanism of retrograde Hot Jupiters. There we % showed the importance of this effect in a verity of planetary and % stellar triple systems. For some examples see Figures % \\ref{fig:hypoplanet3}--\\ref{fig:hypostar}, and \\citet{Naoz10} where we % specifically discussed the evolution of two planet systems, triple % stars and asteroids due to gravitational perturbations from Jupiter. % We also compared our derivation with direct N-body integration and % illustrated the same qualitative evolution.  We also emphasized the % importance of higher-orders approximations, where $\\epsilon_M$ is % significant."
    },
    "1107/1107.2122_arXiv.txt": {
        "abstract": "We describe a search for dust created in collisions between the Saturnian irregular satellites using archival \\emph{Spitzer} MIPS observations. Although we detected a degree scale Saturn-centric excess that might be attributed to an irregular satellite dust cloud, we attribute it to the far-field wings of the PSF due to nearby Saturn. The Spitzer PSF is poorly characterised at such radial distances, and we expect PSF characterisation to be the main issue for future observations that aim to detect such dust. The observations place an upper limit on the level of dust in the outer reaches of the Saturnian system, and constrain how the size distribution extrapolates from the smallest known (few km) size irregulars down to micron-size dust. Because the size distribution is indicative of the strength properties of irregulars, we show how our derived upper limit implies irregular satellite strengths more akin to comets than asteroids. This conclusion is consistent with their presumed capture from the outer regions of the Solar System. ",
        "introduction": "The Solar System's irregular satellites are long thought to have undergone collisions. When only eight were known at Jupiter, \\citet{1981Icar...48...39K} showed that the collision time for objects in the prograde group was less than the Solar System's age. More recently, the discovery of collisional families shows that collisions occurred in the past \\citep{2003AJ....126..398N}. Further evidence lies with the relatively flat size distributions of large irregulars, which \\citet{2010AJ....139..994B} show can be produced from initially steeper distributions by billions of years of collisional evolution. Their collisional evolution is thought to have begun when the irregulars were captured from cold icy regions of the Solar System \\citep[e.g.][]{2007AJ....133.1962N}.  Though the \\citeauthor{2010AJ....139..994B} results are based on reproducing the size distribution of the known irregulars, the destruction of the largest objects must produce vast numbers of fragments. The fragments collide with each other, producing yet more fragments and so on down to dust. The known irregulars are just the largest objects in a continuous size distribution that extends down to the smallest grains that can survive on circumplanetary orbits. It is therefore inevitable that dust associated with irregular satellites exists at some level around the giant planets.  We recently made predictions of the level of this dust based on a collisional model and a simple prescription for the size distribution \\citep{2011MNRAS.412.2137K}. However, these predictions are uncertain for several reasons. The many orders of magnitude between the 10--100km size of known objects and the smallest dust grains means that small differences in the slope of the assumed size distribution result in large differences in the level of dust. This slope depends on the strength law, which is uncertain. There may also be additional loss mechanisms that flatten the size distribution relative to theoretical predictions for a collisional cascade.  \\begin{figure} \\begin{center} \\hspace{-0.35in} \\psfig{width=0.47\\textwidth,figure=saturnpic.eps} \\caption{Simple model of Saturn's irregular satellite dust cloud in Saturnian and sky coordinates (y=0 is the ecliptic). The scale is a linear stretch. Saturn's position is marked by a dot (brightness not to scale). The boxes show the approximate position of the two Phoebe ring (ON,OFF) and two Iapetus (East,West) scans.}\\label{fig:saturnpic} \\end{center} \\end{figure}  Detection of this dust is therefore important for several reasons. Primarily it would provide the strongest evidence yet that the known irregular satellites are just the tip of a continuous size distribution. Because the collisional lifetime of the dust is of order tens of millions of years, any existing dust must be continually replenished by destruction of larger objects, which implies ongoing destruction of objects of all sizes. Characterising the dust population is also important for understanding irregulars themselves, because the slope of the size distribution depends on their strength properties \\citep{2003Icar..164..334O}. More generally, the irregular satellites provide a rare chance to observe how the size distribution in a collisional cascade extrapolates from large objects down to dust, thus informing models used to model extrasolar debris disks.  Additional motivation comes from the existence of dust related to a single Saturnian irregular satellite; the Phoebe ring \\citep{2009Natur.461.1098V}. Impacts that launch grains from Phoebe are proposed as a way to feed the ring. Whether the impactor population is interplanetary, or belongs to a cloud of circumplanetary dust (i.e. due to irregulars) is not known. Detection of a dust cloud would provide compelling evidence in favour of the latter.  A simple way to estimate the spatial distribution of such a grain population is to assume that they follow the orbits of the observed irregulars, which yields a model for the surface brightness distribution of the dust cloud. Figure \\ref{fig:saturnpic} shows this model for Saturn, made by generating a population of particles with random semi-major axes, eccentricities, and inclinations in the range derived for captured Saturn-centric orbits by \\citet{2007AJ....133.1962N}. The scale shows that the cloud extends over a degree in each direction. The smaller vertical extent is the result of a lack of near-polar orbits, which are unstable due to Solar perturbations \\citep{2002Icar..158..434C,2003AJ....126..398N}. The peak surface brightness predicted for Figure \\ref{fig:saturnpic} is 1.25MJy/sr \\citep{2011MNRAS.412.2137K}, but as noted above is uncertain.  Our aim here is to test this prediction for Saturn's dust cloud using the 24$\\mu$m \\emph{Spitzer Space Telescope} observations used to discover the Phoebe ring. The regions of sky covered by these observations relative to Saturn is also shown in Figure \\ref{fig:saturnpic}. The data cover a sufficiently wide region around Saturn to appear promising for a deeper look for a cloud of dust originating from the irregular satellites. In the next sections, we describe our efforts to extract a signal that can be compared with Figure \\ref{fig:saturnpic}. ",
        "conclusions": "\\label{sec:disc}  There are two main points worth discussing; what the issues will be for future observations that attempt to look for the same signal and whether they can be overcome, and what we learn about irregular satellites from our upper limit on the surface brightness of irregular satellite dust.  \\subsection{Constraints on Irregular Satellites}  Although our characterisation of the Spitzer PSF suggests that the excess emission seen in Figure \\ref{fig:sub} is due to the wings of the PSF rather than a dust cloud, the derived upper limit on dust sets interesting constraints on the size distribution of irregular satellites.  The cloud model shown in Figure \\ref{fig:sub} has a total flux of 45Jy. This flux is derived by scaling the model in Figure \\ref{fig:saturnpic} so that a profile at the ON image position matches the level in Figure \\ref{fig:sub}. To estimate the surface area $\\sigma_{\\rm tot}$ in dust required to generate this level of thermal emission, we assume blackbody grains at a temperature of $T= 88$K. Using the relation $F_\\nu = B_\\nu(\\lambda,T) \\, \\sigma_{\\rm tot} / a_{\\rm Sat}^2$ (where $a_{\\rm Sat}=9.5$AU is Saturn's semi-major axis), we therefore find an upper limit to the surface area in dust of $\\sigma_{\\rm tot} = 1.3 \\times 10^{-9}$AU$^2$ (or $2.9 \\times 10^7$km$^2$). For comparison, the projected area of the model in Figure \\ref{fig:saturnpic} is about 0.15AU$^2$ so the limit on the optical depth is extremely low.  This limit assumes grains absorb and emit like blackbodies, but for more realistic grain properties the surface density limit can only be more constraining. Real grains emit inefficiently at wavelengths longer than their physical size, and are therefore hotter than the equilibrium blackbody temperature. The difference here is minor because the smallest grains are relatively large (16$\\mu$m). Depending on the grain properties (porosity, composition, crystallinity) the emission can be the same or a factor few larger. For the highly porous grains inferred for comet-like dust \\citep[e.g.][]{1998A&A...331..291L}, the real grain emission is about a factor two larger than for a blackbody. Because there are uncertainties in the grain properties and derived spectra we retain the more conservative blackbody limit, and note the differences real grains make below.  This upper limit can be converted to a limit on the average size distribution between the smallest grains and largest objects, assuming a single phase size distribution that follows $n(D) = K \\, D^{2-3q}$ \\citep[e.g.][]{2007ApJ...663..365W}, where $D$ is the planetesimal diameter in km, and $K$ the normalisation. We assume the smallest grains are of size $D_{\\rm min} = 16\\mu$m, approximately the smallest grains that can survive in orbit around Saturn \\citep{1979Icar...40....1B,2011MNRAS.412.2137K}. There are about 10 Saturnian irregular satellites larger than 10km in diameter, so the normalisation using the cumulative number is $K = 10 (3q-3)/(10^{3-3q})$. Using the conversion from the size distribution to total surface area $\\sigma_{\\rm tot}$ in grains \\citep{2007ApJ...663..365W} \\begin{equation} \\sigma_{tot} = 3.5 \\times 10^{-17} K (10^{-9} D_{\\rm min})^{5-3q} / (3q-5) \\end{equation} we can solve for the unknown $q$, which yields $q<1.81$. That is, if $q$ were larger, the size distribution would be steeper and there would be more dust than our upper limit. This maximum average size distribution index is similar to the commonly used canonical value of $q=11/6$, derived for a collisional cascade size distribution where strength is independent of size \\citep{1969JGR....74.2531D}. The upper limit on $\\sigma_{\\rm tot}$ is three times lower than would have been predicted with $q=11/6$.  \\begin{figure} \\begin{center} \\hspace{-0.35in} \\psfig{width=0.52\\textwidth,figure=sizedist.eps} \\caption{Cumulative surface area plot showing theoretical distributions and the known irregular satellites of Saturn. The black dotted line is the upper limit on the surface area implied by the MIPS observations. The grey dotted line is the less conservative upper limit.}\\label{fig:sizedist} \\end{center} \\end{figure}  Our predictions in \\citet{2011MNRAS.412.2137K} were based on a two phase size distribution, which fit the size distribution of known irregulars. The two phases arise because object strength sets the size distribution slope \\citep{2003Icar..164..334O}, and there are two different strength regimes \\citep[e.g.][]{1998Icar..135..431D,1999Icar..142....5B,2009ApJ...691L.133S}. Above the $\\sim$0.1km transition size, objects gaining strength from self-gravity have a shallower slope ($q_{\\rm grav}$) than smaller objects, whose strength is derived from material properties (and have a slope $q_{\\rm str}$). The slope in the strength regime is generally negative because larger objects are more likely to have a significant flaw \\citep[e.g.][]{1994Icar..107...98B}. The slope $q_{\\rm str}$ is flatter (less negative) for more porous objects \\citep[e.g.][]{1990Icar...84..226H,2011Icar..211..856H}.  While studies generally agree that the strength and gravity regimes exist, the transition size and dependence of strength on size varies between authors. In predicting a level of 320Jy for Saturn's dust cloud \\citet{2011MNRAS.412.2137K}, we used a transition size of 0.1km, $q_{\\rm str}=1.9$, and $q_{\\rm grav}=1.7$.  Figure \\ref{fig:sizedist} compares our upper limit with the prediction and two other size distributions. All size distributions end at the assumed minimum grain size $D_{\\rm min}=16\\mu$m. The black solid line shows our prediction and the dashed line shows the same distribution with a smaller transition size of 5m. The solid grey line shows the same distribution with a slightly flatter size distribution in the strength regime ($q_{\\rm str}=1.86$). The observed Saturn irregulars sit at the lower right, illustrating the many orders of magnitude in size between the largest and smallest objects, and why predictions are uncertain. We also include a less conservative (but not unreasonable) limit on surface density, assuming that the Iapetus scans contain no dust and that the dust is made of ``real'' comet-like grains (see above), which is a factor four lower.  Figure \\ref{fig:sizedist} shows that there are several ways the surface area can be lower than the upper limit. The surface area decreases as the minimum grain size $D_{\\rm min}$ increases, and for our original prediction would need to be about a factor of ten higher ($\\sim$160$\\mu$m) to be consistent with the upper limit. Models that explore circumplanetary grain dynamics in this context are therefore needed. If the effective minimum grain size is larger our upper limit becomes less constraining, but if grains are typically smaller it is more constraining.  Alternatively, the transition between the strength and gravity regimes may be smaller than the assumed 0.1km. For the same strength dependence, the transition size needs to be 20 times smaller (5m) to satisfy our upper limit (dashed line). Though there are modest variations, no models predict the transition size to be this small \\citep[e.g.][]{1999Icar..142....5B}.  A more likely possibility is that the size distribution in the strength regime is flatter, as illustrated by the grey line. The slope of $q_{\\rm str}=1.86$ corresponds to a strength dependence of $\\propto D^{-0.17}$. For the less conservative limit $q_{\\rm str} = 1.83$ and the strength dependence is $\\propto D^{0.06}$ (i.e. slightly positive). This dependence is flatter than most theoretical models, which lie in the range -0.2 to -0.6 \\citep[e.g.][]{1990Icar...83..156D,1999Icar..142....5B,2009ApJ...691L.133S}. It is significantly flatter than the strength dependence inferred for small asteroids impacting Gaspra and Ida \\citep[-1, ][]{1994Icar..107...84G,1996Icar..120..106G}. A likely reason the strength dependence is flatter than typically predicted or observed is that small irregulars are more porous than asteroids \\citep[e.g.][]{1990Icar...84..226H,2011Icar..211..856H}, and therefore more comet-like \\citep[e.g.][]{1990ApJ...361..260G,2002aste.conf..485B}. Strength properties more akin to comets than asteroids are consistent with the proposed scattering and subsequent capture or irregulars from the outer reaches of the Solar System \\citep[e.g.][]{2007AJ....133.1962N}. The difference between out upper limit of 45Jy and the prediction of 320Jy may therefore be because the strength of small irregular satellites depends less strongly on their size than we assumed, and the size distribution consequently flatter.  However, the strength prescription could be correct because any additional loss processes (other than collisional evolution) would lead to an overall flatter size distribution. The slightly positive strength dependence for our less conservative upper limit suggests that other loss processes do indeed occur. For Saturn, we estimated that Poynting-Robertson (PR) drag is roughly as important as collisional evolution for removing the smallest grains, and therefore possibly an important loss process that may modify the small end of the size distribution \\citep{2011MNRAS.412.2137K,2011arXiv1103.5499W}. For the originally assumed size distribution in Figure \\ref{fig:sizedist}, the surface area would need to turn over around 160$\\mu$m, which is significantly larger than the minimum grain size and therefore unlikely. However, as for the minimum grain size, grain dynamics need to be modelled in detail to study the effect of radiation forces, including other possible effects (e.g. Yarkovsky force).  Observationally, pushing the minimum known sizes of irregular satellites down will better characterise their size distribution over a wider range. However, progress will be difficult. Table 3 in \\citet{2007ARA&A..45..261J} shows that the faintest known irregulars at Saturn have R magnitudes of 24.5,\\footnote{The faintest published survey has $R=24.5$ \\citep{2001Natur.412..163G}, but \\citet{2007ARA&A..45..261J} list an unpublished survey with a limiting magnitude of $R=26$.} and the deepest survey for Uranian irregulars had a 50\\% detection efficiency at $R=26$mag \\citep{2005AJ....129..518S}. Such an improvement for Saturnian satellites would push the smallest objects down by a factor of about 4, to about 1km.  Complementary methods that probe smaller sizes are therefore also needed, such as crater counts. While craters have been counted on Saturnian satellites \\citep[e.g.][]{1982Sci...215..504S,2010Icar..206..485K}, to back out a size distribution of irregulars is complex. For example, because small and large impactors alter the target's surface in different ways, the observed crater size distribution does not necessarily reflect the impactor distribution \\citep[e.g.][]{1994Icar..107...84G,1996Icar..120..106G,2009Icar..204..697R}. Few studies consider irregular satellites as a source of impactors \\citep[see][]{2010AJ....139..994B}, probably because most irregulars were only discovered in the last ten years or so and it was only recently proposed that the known irregulars are the tip of an iceberg of collisional fragments \\citep{2010AJ....139..994B,2011MNRAS.412.2137K}. \\citet{2010AJ....139..994B} make a simple comparison based on their irregular satellite evolution model, but also acknowledge the complications of trying to derive a crater population from a population of irregulars whose size distribution and dynamics evolve as they grind down. Given that the number of irregulars was likely much greater in the past, they may have dominated the impactor population and should be considered when modelling crater counts.  \\subsection{Future observations}  In trying to extract the signal of our proposed dust cloud, we had to consider a number of possible confounding issues. For the ON and OFF scans the Zodiacal contribution could largely be accounted for because the Zodiacal contribution is expected to change little with small changes in elongation. An alternative way to account for ecliptic dust it is to take two observations separated by exactly one year. The first includes the region around the planet, and the second observes the same patch of sky (the planet having moved on in its orbit), and thus has the same line of sight through the ecliptic and the same (extra)galactic background. However, the limiting factor for this study was in fact Spitzer's PSF. PSF characterisation may be the most important issue for future observations of this type.  Given that instrumental PSFs may be hard to avoid and characterise, an additional way to maximise the chance of a detection is to look for dust around a different planet. The argument for doing so could be due to either a better chance of detection (assuming a similar dust cloud is present), or that the dust cloud is likely to be brighter. For the former case, the giant planets are near their equilibrium temperatures so their spectra are roughly similar to that expected for dust. Therefore, a contrast advantage can be gained by looking at Uranus, which is much fainter and whose dust cloud is expected be of similar surface brightness to that predicted for Saturn \\citep{2011MNRAS.412.2137K}. The question of which planet is more likely to harbour detectable dust is uncertain. Uranus may again be more favourable than Saturn because the dust may be less affected by PR drag \\citep{2011MNRAS.412.2137K}. Dust from irregular satellites at Neptune is probably at a lower level due to the proposed disruption of the population by Triton and/or Nereid \\citep{2003AJ....126..398N,2005ApJ...626L.113C}.  While our model considers a large uniform dust cloud, there is also the possibility of smaller scale structures due to individual objects such as the Phoebe ring \\citep{2009Natur.461.1098V}. Because they are non-axisymmetric, these structures should be easier to detect. The best candidates appear to be Phoebe, Nereid, and perhaps the Carme collisional family, because their low inclinations relative to our line of sight increase the surface brightness of any associated dust.\\footnote{A face-on Phoebe ring would be about 15 times fainter or about 0.03MJy/sr. Such a faint ring would not have been detected in the Spitzer observations.}"
    },
    "1107/1107.2408_arXiv.txt": {
        "abstract": "Excesses on positron and electron fluxes measured by ATIC, and the PAMELA and Fermi--LAT telescopes can be explained by dark matter annihilation in our Galaxy. However, this requires large boosts on the dark matter annihilation rate. There are many possible enhancement mechanisms, such as the Sommerfeld effect or the existence of dark matter clumps in our halo. If enhancements on the dark matter annihilation cross section are taking place, the dark matter annihilation in the core of the Earth should also be enhanced. Here we use recent results from the IceCube 40--string configuration to probe generic enhancement scenarios. We present results as a function of the dark matter--proton interaction cross section, $\\sigma_{\\chi p}$ weighted by the branching fraction into neutrinos, $f_{\\nu\\overline{\\nu}}$, as a function of a generic boost factor, $B_F$, which parametrizes the expected enhancement of the annihilation rate. We find that dark matter models which require annihilation enhancements of $\\mathcal{O}(100)$ or more and that annihilate significantly into neutrinos are excluded as the explanation for these excesses. We also determine the boost range that can be probed by the full IceCube telescope. ",
        "introduction": "Introduction}  Several recent observations in our galaxy have found an excess on the electron and/or positron fluxes~\\cite{atic,pamela,fermilat}. However the origin of these events is yet not clear. If they are produced accordingly to standard physics, pulsars~\\cite{puls} might be their sources as well as they could be secondary particles created by shock accelerated hadrons~\\cite{blasi}. Another exciting possibility is that they can be a signature of new physics. It was shown that dark matter annihilation in the galaxy can describe the data~\\cite{bmod}. However most of these models require an enhancement on its annihilation rate. An enhancement mechanism can be found from a variety of possible phenomena, such as the Sommerfeld effect, the existence of dark matter substructures in the galactic halo, or a combination of both~\\cite{cholis2, lattanzi}. \\par  Constraints on a boost factor on the dark matter annihilation rate have been recently derived from the Fermi--LAT diffuse gamma--ray measurement~\\cite{fermicq} and from their analysis of Milky Way satellite galaxies~\\cite{fermibnds}. An independent and neat analysis has been performed by the authors of Ref~\\cite{cmbbnds}, using the fact that dark matter annihilations at recombination time, redshift $\\sim$ 1000, would have injected secondary particles that would have affected the recombination processes. The measured power spectrum of the cosmic microwave background (CMB) can thus be used to set limits on the strength of the dark matter self--annihilation cross section. The IceCube collaboration has also performed an analysis searching for dark matter annihilations in the galactic halo by searching from an excess diffuse neutrino flux over the expected atmospheric flux~\\cite{ic22halo}. The results can also be used to set limits on a boost in the annihilation cross section.\\par  These methods show some degree of complementarity since each one provides better sensitivity to a different range of the mass of the dark matter particles or to different annihilation channels. They also rest on different assumptions and approximations. Fermi and the CMB method are competitive in the low mass range (dark matter massed $\\lesssim$ 1~TeV), while IceCube reaches 10 TeV, and can also probe annihilation directly into neutrinos. There is some halo-model dependency when setting limits on the annihilation cross section from the observations of satellite galaxies: the expected signal depends on the degree of cuspiness of the assumed halo. On the other hand, the CMB analysis does not depend  on the shape of dark matter haloes, since there are no gravitationally bound structures at the redshift considered, but it depends on assumptions about the fraction of the energy released by the annihilating dark matter particles and how it is absorbed by the surrounding medium. \\par  In this paper we follow the alternative approach proposed in Ref.~\\cite{paddy}. The authors argue that an enhancement of the dark matter annihilation rate should also boost the neutrino flux from dark matter annihilations in the center of the Earth. A thermal annihilation cross section implies that dark matter capture and annihilation has not yet reached equilibrium in the Earth. However it is not necessary that the annihilation cross section has to remain the same as during the freeze--out period. If the annihilation rate is somehow enhanced in the post freeze--out period, the equilibrium might already have been achieved. As it will be shown in the next Section, the annihilation rate reaches its maximum at the equilibrium, and then depends only in the capture rate. A boost in the annihilation rate increases the flux of annihilation products until it reaches its maximum. In this case the neutrino flux from the center of the Earth will be large enough to be detected by telescopes such as IceCube. In short, if the dark matter annihilation rate is enhanced, the timescale for equilibrium diminishes and the flux of annihilation products is expected to be much larger than the one away from equilibrium.\\par  We use the recently published IceCube results on a search for a diffuse flux of muon neutrinos~\\cite{ic40diff} to set limits on a boost factor in the dark matter annihilation cross section. IceCube has measured neutrinos coming from near or below the horizon in the energy range 332~GeV and 84~TeV using data taken with the 40--string detector configuration (IceCube--40). The analysis includes all events coming from near or below the horizon, and the result is compatible with the expected atmospheric neutrino flux (see also Ref.~\\cite{ic40atm}). It is also generic enough to allow comparison with our prediction of the flux of muon neutrinos produced in dark matter annihilations in the center of the Earth. We determine this flux by simulating the annihilation of WIMP--type particles in the center of the Earth and propagating the neutrinos to the detector. A significant neutrino flux from dark matter annihilation should be seen above the expected atmospheric neutrino flux. If not, limits can be set on the model used to determine the dark matter signal. Our analysis shows that models which require very large boosts on the dark matter annihilation rate in order to explain the excess seen in the galactic positron and electron flux, and have an annihilation channel into neutrinos, are ruled out.  \\par    In the next Section we describe dark matter capture and annihilation in the Earth. In Sections~\\ref{sec:neutrinos} and~\\ref{sec:constrains} we describe the signal simulation and calculate the expected number of events in the IceCube--40 detector. We then compare our results with the IceCube--40 published results, showing the boost factor range that is excluded. Finally, in Section~\\ref{sec:prediction} we estimate the sensitivity region for the full 86--string detector. ",
        "conclusions": "Conclusions}  If dark matter annihilation in our galaxy is responsible for the excesses seen by ATIC~\\cite{atic}, PAMELA~\\cite{pamela} and Fermi--LAT~\\cite{fermilat}, an enhancement on the dark matter annihilation rate in the Earth should also be foreseen. This enhancement might bring the equilibrium among the dark matter capture and annihilation rates to occur much earlier than the timescale expected from a  purely thermal annihilation rate. In this case the neutrino flux from dark matter captured in the Earth can be quite large. We present results as a function of the  dark matter--proton interaction cross section, $\\sigma_{\\chi p}$ weighted by the branching fraction into neutrinos, $f_{\\nu\\overline{\\nu}}$, as a function of a generic boost factor, $B_F$, which parametrizes the expected enhancement of the annihilation rate. In this sense, it is important to note that our results do not depend on the details of the mechanism which enhances the annihilation cross section. We have used two benchmark models, a 500~GeV and 1~TeV generic WIMP annihilating in the center of the Earth to scan the  ($\\sigma_{\\chi p} \\times f_{\\nu\\overline{\\nu}}$, $B_F$) parameter space and set constrains on this 2-dimensional space using current IceCube results. Our calculations assume that the dark matter velocity distribution is Gaussian as well as that dark matter collected in the Earth's core follows an isothermal distribution. \\par  In order to explain the positron and electron excesses that are seen, models also have to cope with the fact that the antiproton spectrum measured by PAMELA~\\cite{pamantip} is in full agreement with the expectation from secondary production of antiprotons from propagation of cosmic rays in the galaxy. This rules out as an explanation of the excess many dark matter models with preferred annihilation into heavy products, producing many antiprotons. Therefore leptophilic models, which propose dark matter annihilation exclusively into leptons, are favored as an explanation to the excesses found. We have shown that, when 500 (1000)~GeV dark matter annihilates into a large fraction of neutrinos, annihilation boost factors of the $\\mathcal{O}(100)$ and above are already constrained by our analysis at a 5$\\sigma$ level, or higher, depending on the interaction cross section assumed. Thus, leptophilic models~\\cite{leptop} which favor primary neutrino production are constrained by our results. \\par  In order to compare this analysis to others, we show in Figure~\\ref{fig:cmp} limits on $\\left< \\sigma_A \\times v \\right>$ as a function of $m_\\chi$ from Fermi~\\cite{fermibnds}, CMB~\\cite{cmbbnds} and IceCube~\\cite{icebnds}. Our bounds as a function of $\\left< \\sigma_A \\times v \\right>$, can be determined from equation~\\ref{eq:rann}. It can also easily be visualized in Figure~\\ref{fig:cl}, where a $3 \\sigma$ significance is shown as the edge of the shaded area: the limit on the annihilation cross section is determined from each boost factor value on this curve. In Figure~\\ref{fig:cmp} these are shown for two choices of $\\sigma_{\\chi p}$, $3 \\times 10^{-44}$~cm$^2$ (blue stars) and $1 \\times 10^{-44}$~cm$^2$ (blue solid squares), which are below the current Xenon limit~\\cite{xe100}.  \\begin{figure}[t] \\vspace*{-1.5cm} \\hspace*{-.5cm} \\includegraphics[width=4.15in,height=5.25in]{cmp_4.pdf} \\vspace*{-2.5cm} \\caption{Bounds on the dark matter annihilation cross section $\\sigma_{A}$ versus mass $m_\\chi$ by Fermi~\\cite{fermibnds}, IceCube~\\cite{icebnds}, analysis over CMB data~\\cite{cmbbnds} and this work, where two different $\\sigma_{\\chi p}$ values of $3 \\times 10^{-44}$~cm$^2$ (blue star) and $1 \\times 10^{-44}$~cm$^2$ (blue solid square) are assumed. Here we also plot our results assuming a 5 TeV dark matter mass. Our limits correspond to a $3 \\sigma$ significance level, while the others are at 90\\% CL.} \\label{fig:cmp} \\end{figure}  Our results and IceCube's, are the only ones to probe annihilation into neutrinos. Other limits coming from searches with gammas from satellite galaxies by Fermi~\\cite{fermibnds} or from the analysis based on the CMB~\\cite{cmbbnds} imprint at high redshifts, are complementary, since not only they probe different annihilation channels, but also rest on different underlying assumptions. \\par  We also investigated the reach of the completed IceCube 86-string detector, and we present results of its sensitivity in the  ($\\sigma_{\\chi p} \\times f_{\\nu\\overline{\\nu}}$, $B_F$) parameter space. We have estimated how using neutrino energy information could improve the analysis, and in Fig.~\\ref{fig:cl} we show the expected sensitivity for three different assumptions of the detector energy resolution."
    },
    "1107/1107.3638_arXiv.txt": {
        "abstract": "{A complete periodic star extraction and classification scheme is set up and tested with the Hipparcos catalogue. The efficiency of each step is derived by comparing the results with prior knowledge coming from the catalogue or from the literature. A combination of two variability criteria is applied in the first step to select 17\\,006 variability candidates from a complete sample of 115\\,152 stars.  Our candidate sample turns out to include 10\\,406 known variables (i.e., 90\\% of the total of 11\\,597) and 6600 contaminating constant stars. A random forest classification is used in the second step to extract 1881 (82\\%) of the known periodic objects while removing entirely constant stars from the sample and limiting the contamination of non-periodic variables to 152 stars (7.5\\%).  The confusion introduced by these 152 non-periodic variables is evaluated in the third step using the results of the Hipparcos {\\em periodic} star classification presented in a previous study (Dubath et al.~\\cite{paper1}).} ",
        "introduction": "\\label{sec:intro}   Current and forthcoming photometric surveys are monitoring very large numbers of astronomical targets providing a fantastic ocean for fishing interesting variable objects. However, because of the large numbers involved, their extraction requires the use of fully automated and efficient data mining techniques. In this contribution, we use the Hipparcos data set to investigate the performance of a complete and automated scheme for the identification and the classification of periodic variables. As shown in Fig. \\ref{fig:1}, we study a three step process. In the first step variable candidates are separated from the objects most likely to be constant. This saves significant processing time as period search is performed only in the subset of variable candidates. The validity of the detected periods is established in the second step, which separates truly periodic from non-periodic objects. The third step is the classification of periodic variables into a list of types (only a sub-set of them are shown in Fig. \\ref{fig:1}). This step is presented in details in Dubath et al.~\\cite{paper1}.  To avoid unnecessary repetition, this first paper is referred to for a full description of the classification attribute calculation and of the details of the random forest methodology.   \\begin{figure}[h] \\sidecaption \\includegraphics[width=\\columnwidth]{dubath_fig1.pdf} \\caption{Illustration of the steps used in this study to identify and classify variable sources} \\label{fig:1}       % \\end{figure}  This three step organisation represents a particular option. Alternatives are also being considered, but they are outside the scope of this contribution as is the classification of non-periodic variables which is the subject of another study (Rimoldini et al., in preparation). ",
        "conclusions": ""
    },
    "1107/1107.4368_arXiv.txt": {
        "abstract": "We present evidence that the incidence of active galactic nuclei (AGNs) and the distribution of their accretion rates do not depend on the stellar masses of their host galaxies, contrary to previous studies. We use hard ($2-10$ keV) X-ray data from three extragalactic fields (XMM-LSS, COSMOS and ELAIS-S1) with redshifts from the Prism Multi-object Survey to identify \\amend{242} AGNs with $L_\\mathrm{2-10\\; keV}=10^{42-44}$ erg s$^{-1}$ within a parent sample of $\\sim$25,000 galaxies at $0.2<z<1.0$ over $\\sim3.4$~deg$^{2}$ and to $i\\sim23$.  We find that although the fraction of galaxies hosting an AGN at fixed X-ray luminosity rises strongly with stellar mass, the \\emph{distribution} of X-ray luminosities is independent of mass.  Furthermore, we show that the probability that a galaxy will host an AGN can be defined by a universal Eddington ratio distribution that is \\emph{independent} of the host galaxy stellar mass and has a power-law shape with slope $-$\\amend{0.65}.  These results demonstrate that AGNs are prevalent at all stellar masses in the range $9.5<\\log \\mathcal{M}_*/\\mathcal{M}_\\odot<12$ and that the same physical processes regulate AGN activity in all galaxies in this stellar mass range. While a higher AGN fraction may be observed in massive galaxies, this is a selection effect related to the underlying Eddington ratio distribution.  We also find that the AGN fraction drops rapidly between $z\\sim1$ and the present day and is moderately enhanced (factor$\\sim2$) in galaxies with blue or green optical colors.  Consequently, while AGN activity and star formation appear to be globally correlated, we do not find evidence that the presence of an AGN is related to the quenching of star formation or the color transformation of galaxies. ",
        "introduction": "\\label{sec:intro}  Supermassive black holes (SMBHs) are now thought to reside at the centers of most, if not all, galaxies with central stellar bulges \\citep{Kormendy95,Kormendy01}. These galaxies must have undergone periods of intense accretion activity to build their black holes, during which they would be observed as active galactic nuclei (AGNs). The mass of the SMBH is known to be tightly correlated with both the luminosity \\citep{Magorrian98} and velocity dispersion \\citep{Ferrarese00,Gebhardt00,Tremaine02} of the stellar bulge, indicating that the growth of the SMBH and the processes that build up the stellar mass of the host galaxy must be related, \\amend{although it is unclear whether a causal connection between these processes is required \\citep[see][]{Peng07,Jahnke11}.} In addition, the star formation rate density as a function of redshift is found to correlate with the black hole accretion rate density out to high redshifts \\citep[e.g.,][]{Silverman08,Aird10}, indicating that these processes may be \\amend{controlled} by a common mechanism. Indeed, a number of studies have proposed that AGNs provide feedback that quenches star formation in their host galaxies and is crucial to control the transformation of blue, star-forming galaxies to red, passively evolving systems \\citep[e.g.,][]{Croton06,Hopkins07,Schawinski07}, although there remains a lack of direct observational evidence of this process. Nevertheless, what triggers and regulates AGN activity and how this relates to the co-evolution of AGNs and their host galaxies remain key, open questions in extragalactic astrophysics.  To understand how AGNs and galaxies co-evolve we must study the types of galaxies that tend to host AGNs. Galaxies can be separated into a variety of classifications, many of which have significant overlap. Many studies have shown that the positions of galaxies in the optical color--magnitude diagram are strongly bimodal, falling into two distinct regions associated with passively evolving galaxies on the ``red sequence\" and star-forming galaxies in the ``blue cloud\" \\citep[e.g.,][]{Blanton03,Baldry04,Bell04}; a narrow ``green valley\" separates these two populations. The total stellar mass of a galaxy, \\Mstel, is also a key parameter that tracks the net result of star formation activity throughout a galaxy's lifetime, as well as the hierarchical build up of galaxies through mergers. While both the red and blue galaxy populations span a wide range of stellar masses, red galaxies tend to dominate the population at higher stellar masses, whereas blue galaxies generally have lower stellar masses. Over cosmic time the total combined mass of galaxies on the red sequence is found to increase, indicating that star formation must be quenched in blue galaxies, \\amend{which subsequently allows their stellar populations to evolve passively and transfers the galaxies to the red sequence \\citep{Bell04,Faber07,Brown07}.} Morphology is also closely correlated with galaxy color and stellar mass---massive, red galaxies are generally early-type, bulge-dominated systems, whereas blue galaxies tend to have late-type morphologies with substantial disk components \\citep[e.g.,][]{Blanton09}.  Which of these galaxies host AGNs? \\amend{Given the correlation between the overall AGN accretion and star formation rate densities \\citep{Silverman08,Aird10},} we might expect to find AGNs predominantly \\moreamend{within blue star-forming galaxies, which tend to have low stellar masses.} A number of studies show that this is not the case---the fraction of galaxies hosting AGNs increases in higher stellar mass, bulge-dominated galaxies---although the subject remains under much debate. \\citet{Kauffmann03} found that narrow-line AGNs (identified by their emission-line ratios) are preferentially in massive, bulge-dominated galaxies (\\Mstel$\\gtrsim10^{10}$\\Msun) in the local universe. \\amend{\\citet{Kauffmann03} also found evidence for recent star formation in the host galaxies of the more luminous AGNs, although the host galaxies of the lower luminosity AGNs generally had older stellar populations, similar to most early-type galaxies \\citep[see also][]{Kauffmann07}. } The fraction of galaxies that host radio-loud AGNs is even more strongly dependent on stellar mass,  rising rapidly from $<0.01$\\% at \\Mstel$\\approx 10^{10}$ \\Msun\\ to $>30$\\% at \\Mstel$\\gtrsim 5\\times 10^{11}$\\Msun \\citep[]{Best05}; it has long been established that radio-loud AGNs are primarily found in massive, elliptical galaxies \\citep[e.g.,][]{Matthews64,Dunlop03}. \\citet{Schawinski07} studied the optical colors of elliptical galaxies in the local universe and found that those hosting narrow-line AGNs were generally on the red sequence\\amend{; these} results suggest that star formation and AGN activity do not occur at the same time, which could be explained by the onset of AGN activity  (possibly during an early, highly obscured phase of the AGN lifetime) being responsible for quenching star formation activity within galaxies and transferring them from the blue cloud to the red sequence \\citep[e.g.,][]{Hopkins06b}. \\amend{ However, direct observational evidence of the AGN feedback process is lacking. A detailed study of the prevalence of AGNs as a function of stellar mass and galaxy color is required to shed light on the relationship between the growth of galaxies, the progress of star formation, and the accretion activity of the central SMBH. }  Deep X-ray surveys have enabled efficient identification of obscured and low-luminosity AGNs where the host galaxy dominates the optical light, allowing studies of the host galaxy properties to be pushed out to much higher redshifts \\citep[$z\\lesssim4$;][]{Xue10}. Various studies have shown that the fraction of galaxies hosting X-ray AGNs rises strongly with increasing \\Mstel\\ at all redshifts from $z\\sim0.4-4$ \\citep[e.g.,][]{Alonso-Herrero08,Brusa09b,Xue10}, similar to the optical narrow-line AGNs identified in the local universe \\citep{Kauffmann03}. \\citet{Nandra07b} investigated the positions of X-ray AGN hosts in the optical color--magnitude diagram at $z\\sim1$ and found that X-ray-selected AGNs are predominantly in luminous galaxies but avoid the strongest star-forming galaxies and thus lie in the reddest part of the blue cloud, in the green valley, or on the red sequence. These results indicate that AGN activity may be associated with the transformation of blue galaxies and the quenching of star formation, indicating that AGN feedback could potentially play a key role in this process. Similar results were found by \\citet{Coil09} for a larger sample of X-ray AGNs, and the same pattern appears to hold at lower redshifts \\citep{Hickox09,Georgakakis11}. However, \\citet{Silverman09} and \\citet{Xue10} have demonstrated the importance of stellar-mass selection effects, which may bias prior work. \\citet{Silverman09} found that X-ray AGN activity was associated with ongoing star formation, and thus the AGN fraction is higher in the blue cloud. \\citet{Xue10} found that the AGN fraction rises strongly with increasing \\Mstel, as described above, but when considering samples matched in stellar mass there was no evidence for specific clustering of AGN host galaxies on the color--magnitude diagram. \\citet{Cardamone10b} also studied the relation of AGN host galaxy masses and colors; they found that the colors of AGN host galaxies of a given stellar mass followed the same bimodality as the non-AGN galaxy population, after correcting for the effects of dust. Thus, despite much recent progress, it remains unclear as to whether there is an association between AGN activity and the color transformation of galaxies, and how this relates to the predominance of AGNs in massive galaxies.  Further insight into the processes that trigger AGNs and how these relate to the evolution of galaxies may be possible by directly probing the distribution of AGN activity, the active growth of SMBHs via accretion that produces a luminous AGN. A number of studies have  constrained the shape and evolution of the AGN luminosity function, which traces AGN activity, over a wide redshift range \\citep[e.g.,][]{Ueda03,Ebrero09,Yencho09,Aird10,Assef11}. The luminosity function has a steep double power-law shape with a break at a characteristic luminosity, $L_*$. The rapid decline in the AGN accretion rate density since $z\\sim1$ is a result of a downward shift in $L_*$ for the overall population, but the luminosity function retains \\amend{essentially} the same shape \\citep{Barger05,Aird10,Assef11}. This indicates that the processes responsible for triggering AGNs, fueling black hole growth and controlling the lifetimes and duty cycles of AGNs are essentially the same but are moving towards generally less luminous systems between $z\\sim1$ and the present day. However, few observational studies have attempted to simultaneously study the distribution of AGN activity as traced by the AGN luminosity and the prevalence of AGNs as a function of the host galaxy properties such as color and stellar mass.  Understanding both the predominance of AGNs in massive galaxies and any correlation with host galaxy colors is vital to reveal the role and effect of AGNs on the evolution of galaxies. \\amend{In addition, to constrain the processes that trigger AGNs and drive black hole growth we must track the evolution of AGN accretion rates and the masses of their central SMBHs, although direct measurement of these quantities is extremely difficult, if not impossible for large samples of distant AGNs and galaxies.} \\amend{In this paper, we use the X-ray luminosity as a tracer of the AGN accretion rate and investigate the relationship with the total stellar mass and colors of galaxies. While the relationship between these observed properties and the underlying physical parameters remains uncertain, careful investigation including consideration of observational biases and selection effects is vital to reveal and constrain the connection between the growth of galaxies and their SMBHs. Our study focuses on } intermediate redshifts ($z<1$), a key time period of the universe over which star formation and AGN activity have rapidly declined from the peak at $z\\sim1-2$. We use the PRIsm MUlti-object Survey \\citep[PRIMUS;][Cool et al., in preparation]{Coil10}, a recently completed faint-galaxy survey that used a prism to obtain low-resolution spectra and determine redshifts for $\\sim 120,000$ galaxies over $\\sim 9 $ \\degsq\\ with $z\\sim 0.2-1.2$, providing the largest currently available spectroscopic sample of galaxies at these redshifts. X-ray data are available for $\\gtrsim 3$ \\degsq\\ of PRIMUS from a variety of \\textit{Chandra} and \\textit{XMM-Newton} survey programs. We use these X-ray data to identify AGNs within the PRIMUS sample, enabling us to efficiently select AGNs down to low luminosities and with moderate obscuration. For such sources the optical light probes the host galaxies, and we can compare properties with the extremely large non-AGN galaxy sample provided by PRIMUS. Section \\ref{sec:data} briefly describes our data, providing details on the compilation of X-ray source lists and matching to PRIMUS sources. Section \\ref{sec:stellarmass} outlines our derivation of stellar masses and Section \\ref{sec:sample} describes in detail the selection of our AGN and galaxy samples. We investigate how the probability of a galaxy hosting an AGN is related to the stellar mass of the galaxy and the AGN X-ray luminosity in Section \\ref{sec:frac} for two wide redshift bins. In Section \\ref{sec:zevol} we determine the redshift dependence of the AGN fraction and perform a more robust analysis of the stellar mass and luminosity dependence that accounts for the evolution with redshift via a maximum-likelihood fitting approach. In Section \\ref{sec:eddratio} we show that the probability that a galaxy hosts an AGN is primarily determined by a power-law distribution of Eddington ratios with a slope of $-0.65$ and a normalization that depends on redshift but does \\emph{not} depend on the stellar mass of the host galaxies. We investigate the dependence of the AGN fraction on galaxy color in Section \\ref{sec:col} and find a weak enhancement of AGNs in galaxies with bluer colors. We discuss the implications of our results in Section \\ref{sec:discuss}. Section \\ref{sec:summary} summarizes our results and overall conclusions.  Throughout this paper we adopt a flat cosmology with $\\Omega_\\Lambda=0.7$ and $h=0.7$. All magnitudes are on the AB system. ",
        "conclusions": "\\label{sec:discuss}  The results presented in this paper reveal how the prevalence and distribution of moderately obscured, moderate-luminosity AGN accretion activity is related to the properties of their potential host galaxies during a key epoch in the lifetime of the universe, $0.2<z<1.0$. We find that the distribution of X-ray luminosities for all stellar masses has the same power-law shape. The probability of a galaxy hosting an AGN above a given \\emph{X-ray luminosity} increases strongly with the stellar mass of the galaxy. However, we show that this effect may be primarily driven by the Eddington ratio distribution. The distribution of Eddington ratios has a power-law shape; a low fraction of galaxies ($\\sim 0.1$\\%) host AGNs accreting close to the Eddington limit (\\Lamedd$\\gtrsim0.1$), whereas $\\sim 10$\\% of galaxies contain an SMBH growing at a very low rate (\\Lamedd$\\sim 10^{-4}$). The probability of hosting an AGN of a given Eddington ratio is independent of stellar mass.  Alternatively, we can interpret these results in terms of a specific accretion rate: the rate of black hole growth relative to the stellar mass of the galaxy. The distribution of specific accretion rates is a universal function that is independent of stellar mass (see Section \\ref{sec:specaccrate}). Our conclusions hold regardless of uncertainties in the black hole mass and thus the absolute value of the Eddington ratios for the AGN.  We find that AGNs are \\emph{not} predominantly in massive, red galaxies. More AGNs may be found in red galaxies due to a combination of two selection effects: (1) it is easier to detect the more prevalent, low Eddington ratio (low specific accretion rate) AGNs in galaxies with higher stellar masses as the AGNs will be intrinsically brighter; and (2) more massive galaxies tend to have redder colors.  There is, however, a mild dependence of the prevalence of AGN activity on galaxy color. While AGNs are found in galaxies across the whole range of colors, there appears to be an enhancement (factor $\\sim 2$) in the probability of hosting an AGN for galaxies with \\emph{blue} colors, both within the blue cloud and the green valley. Thus there may be a weak association between AGN activity and star formation within a host galaxy. Nevertheless, a lack of star formation does not preclude AGN activity---the distribution of Eddington ratios in red galaxies follows approximately the same power-law shape as in blue galaxies, spanning the full range of accretion rates.  AGN activity also evolves strongly with redshift. We find that the probability of a galaxy hosting an AGN has decreased rapidly since $z\\sim1$, but the distribution of Eddington ratios retains the same power-law shape. Thus, the same physical process or processes are likely to be responsible for triggering and fueling AGNs throughout this epoch, but they must diminish rapidly over the time period (see Section \\ref{sec:agntriggering}).  In this section we discuss the implications of our results. In Sections \\ref{sec:eddratiohighz} and \\ref{sec:eddratiolocal} we compare our findings with previous studies of the Eddington ratio distribution at similar redshifts to our study and in the local universe respectively. In Section \\ref{sec:agntriggering} we discuss the implications of our results for models of AGN triggering and fueling processes over the last 8 billion years. In Section \\ref{sec:hostcolor} we discuss our findings on the dependence of AGN accretion activity on host-galaxy color and compare with prior studies of color--magnitude relations for AGNs. Section \\ref{sec:role} discusses the role of AGNs in galaxy evolution. Section \\ref{sec:issues} discusses the remaining issues and uncertainties in our work and prospects for future studies.  \\subsection{Comparison with prior studies of the Eddington ratio distribution at $z\\gtrsim0.2$} \\label{sec:eddratiohighz}  Several prior studies have presented distributions of Eddington ratios for X-ray-selected samples of AGNs at $z\\gtrsim0.2$. \\citet{Hickox09} found that X-ray-selected AGNs in AGES at $0.25<z<0.8$ have a wide range of Eddington ratios spanning $10^{-3}<\\lamedd<1$, with a broad, roughly lognormal distribution peaking at $\\lamedd\\approx0.02$. However, the incompleteness effects due to the X-ray sensitivity were not taken into account. \\citet{Simmons11} determined Eddington ratios for X-ray sources in the GOODS fields (using very deep X-ray data from both the \\chandra\\ Deep Fields) at $0.25<z<1.25$ and found that the sample was dominated by objects accreting at very low Eddington ratios ($\\lamedd<0.01$) with a median of $\\lamedd=0.006$; the extremely deep X-ray data mitigate incompleteness effects and enable detection of very low Eddington ratio sources. \\citet{Brusa09b} also presented an Eddington ratio distribution for obscured, moderate-luminosity, X-ray AGNs identified in the \\chandra\\ Deep Field-South at higher redshifts: $1<z<2$ and $2<z<4$. Similar, broad distributions of Eddington ratios were found in both redshift bins with medians $\\lamedd\\sim0.02-0.08$. None of these prior papers have tracked or corrected for X-ray selection effects to reveal the underlying distribution of Eddington ratios.  Recently, \\citet{Trump11} proposed that there are two distinct distributions of Eddington ratios for X-ray-selected AGNs corresponding to two different modes of accretion. They find that broad-line QSOs are predominantly at higher \\Lamedd$\\sim0.1$, consistent with other studies of the Eddington ratio distribution of QSOs \\citep[e.g.,][]{Kollmeier06,Gavignaud08,Kelly10}. They propose that these AGNs are fed by a thin accretion disk with a disk wind containing the broad-line region and some obscuring material. Narrow-line AGNs or X-ray AGNs lacking optical AGN signatures are shown to have lower \\Lamedd$\\sim10^{-3}$ and may be powered by a geometrically thick, radiatively inefficient accretion flow. Our sample was explicitly limited to ``obscured\" AGNs that lack broad lines in their optical spectra (allowing us to probe the host galaxy properties in the optical light) and will thus correspond to the second class proposed by \\citet{Trump11}. Consistent with their findings, we have shown that these AGNs are predominantly at low Eddington ratios. In fact, we find a power-law relation that extends to very low \\Lamedd$\\sim10^{-4}$. Therefore, the physical processes that trigger and fuel low-level AGN accretion activity \\amend{may} also be common and take place in galaxies of all stellar masses.  Our results are consistent with previous studies of Eddington ratio distributions for moderately obscured, moderate-luminosity AGNs at $z\\gtrsim0.2$. However, we show that the true distribution of Eddington ratios is a power law that continues to rise to lower values of \\Lamedd. This result is reliant on careful analysis of X-ray sensitivity effects relative to the stellar masses of the host galaxies. Our directly observed distribution of Eddington ratios has a broad, approximately lognormal distribution with a peak at \\Lamedd$\\sim10^{-2}$ (see gray histogram in final panel of Figure \\ref{fig:leddfrac_massbin}) as found in prior works but this is significantly skewed by X-ray incompleteness effects and does not represent the true \\Lamedd\\ distribution. We note that our study has excluded BLAGNs which may dominate the population at high \\Lamedd\\ and could correspond to a distinct accretion mode.   \\subsection{Comparison with prior studies of the local Eddington ratio distribution} \\label{sec:eddratiolocal}  The lack of an extremely large area, sensitive X-ray survey at $>2$ keV has meant that the majority of studies of the Eddington ratios of AGNs in the local universe ($z\\lesssim0.1$) have relied on large-area optical spectroscopic surveys to identify AGNs from their optical line-emission, in particular the Sloan Digital Sky Survey \\citep[SDSS; e.g.,][although see also \\citealt{Georgakakis11}]{Kauffmann03,Heckman04,Kauffmann04,Hao05,Kewley06,Schawinski07}. As discussed in Section \\ref{sec:intro}, such studies have found that the fraction of galaxies hosting optically selected narrow-emission-line (Type 2) AGNs rises steeply with increasing stellar mass and that Type 2 AGNs are found predominantly in massive galaxies \\citep[$\\mstel\\gtrsim10^{10} \\msun$;][]{Kauffmann03,Best05}.  \\citet[hereafter H04]{Heckman04} studied the present-day growth rates of black holes in SDSS galaxies using the O{\\sc iii} luminosity as a tracer of AGN accretion activity. They found that the distribution of Eddington ratios is highly dependent on black hole mass; while $\\sim0.5$\\% of black holes with $\\mbh=10^{7}\\msun$ are accreting at Eddington ratios $\\lamedd>0.1$, less than 0.01\\% of higher mass black holes ($\\mbh\\approx3\\times10^{8}\\msun$) are accreting at these high Eddington ratios. \\citet[hereafter K09]{Kauffmann09} built on this work and proposed that black hole growth in the local universe is driven by two distinct processes: (1) a self-regulated mode of growth in young galaxies with plentiful cold gas that leads to a broad lognormal distribution of Eddington ratios peaking at \\Lamedd$\\sim10^{-2.5}$, and (2) a passive mode in galaxies with older central stellar populations where the supply of fuel to the black hole is driven by stellar mass loss from  stars, resulting in a power-law distribution of Eddington ratios. This is somewhat different to our results at higher redshifts using X-ray selection but probing similar ranges of Eddington ratios and black hole masses (roughly traced by the galaxy stellar mass). We find that the distribution of Eddington ratios is the same for all stellar masses and is described by a power law; we find no evidence for an additional lognormal component and that the shape of the Eddington ratio distribution is similar for young and old stellar populations (i.e., blue and red galaxies).  It is unclear how to reconcile our findings with the results of H04 and K09 without further study to directly compare X-ray and optical probes of AGN accretion in both the local and high-redshift universe. The discrepancies could be related to biases inherent to the two very different selection techniques that affect both the completeness of the AGN samples and the estimates of the Eddington ratios. For example, AGNs in the lognormal mode identified by K09 could be heavily obscured and thus are not detected in X-rays or are detected but their luminosities (and thus Eddington ratios) are significantly underestimated. Alternatively, it is possible that the O{\\sc iii} emission line may not provide an accurate tracer of the black hole accretion at low Eddington ratios or in galaxies with current star formation.  Another possibility is that the differences between our results and those of H04 and K09 are due to evolution with redshift. Indeed, H04 consider volume-averaged growth rates and show that the most massive black holes in their sample are \\emph{not} growing at significant rates at the present day and \\emph{must} have undergone periods of rapid growth in the earlier universe. At higher redshifts we may be seeing this earlier phase of relatively rapid growth for all galaxies over the range of stellar masses (and thus black hole masses) that we are able to probe, although our results do not show any direct evidence for such ``downsizing'' behavior (see further discussion in Section \\ref{sec:agntriggering} below).  The relative importance of the two accretion modes proposed by K09 could also change with redshift. We find that the distribution of Eddington ratios has a power-law shape with a normalization that increases rapidly with redshift as $\\sim(1+z)^{3.5}$. The lognormal accretion mode seen by K09 may not evolve or may evolve only weakly with redshift and thus would not be seen in our sample. This may, however, conflict with the K09 interpretation that the accretion is driven by stochastic accretion of cold gas, as the supply of cold gas should be more plentiful in the high-redshift universe.  Further work and a direct comparison of X-ray and optical probes of AGN accretion modes in both the local and high-redshift universe is required to resolve these issues.   \\subsection{The Evolution of AGN Triggering and Fueling} \\label{sec:agntriggering}  Our observed AGN Eddington ratio distribution can provide key insights into the physical processes that are responsible for triggering and fueling AGNs at $z\\sim0.2-1$. One of our major findings is that the probability of hosting an AGN for galaxies throughout our stellar mass range ($9.5<\\log\\mstel/\\msun<12$) is primarily determined by the Eddington ratio distribution and is independent of the stellar mass. Not only is the shape of the distribution of accretion rates (tracked by the Eddington ratios) the same, the overall normalization does not depend on the stellar mass. Thus, AGN activity may take place in any galaxy, regardless of the stellar mass. The universality in both the prevalence of black hole growth and the distribution of accretion rates (Eddington ratios) indicates that the same underlying physical processes are responsible for triggering and fueling AGN activity in all moderately massive galaxies.  We also find that the fraction of galaxies hosting AGNs above a given Eddington ratio evolves strongly with redshift with the form $(1+z)^{3.5\\pm0.5}$ to $z=1$. The shape of the distribution remains the same, indicating that the same processes are regulating AGN activity but are undergoing a rapid evolution with redshift. This evolution could take two possible forms: (1) the fraction of accreting systems is higher at higher redshift, or (2) the characteristic Eddington ratio is higher at higher redshift. In both scenarios, a given physical process must trigger AGN activity across our redshift range, resulting in the same distribution of Eddington ratios with time. However, in the first scenario the fraction of galaxies with this process is higher at earlier times, while in the second scenario the AGNs are more rapidly accreting at earlier times.  We cannot distinguish between these two scenarios solely with the results presented here.  However, studies of the evolution of the X-ray luminosity function (XLF) of AGNs, which traces the overall distribution of AGN accretion activity, find that at low redshifts ($z\\lesssim1$) the evolution is dominated by luminosity evolution \\citep{Barger05,Aird10}. The characteristic luminosity, $L_*$, corresponding to the break in the double-power-law form of the XLF, shifts to lower luminosities between $z\\sim1$ and the present day. The Eddington ratio distribution presented here will track the evolution of the faint end of the XLF. Indeed, if we translate the luminosity evolution of $L_*$ into an effective change in normalization at the faint end of the XLF we find evolution of the form $(1+z)^{3.9\\pm0.3}$ \\citep[based on the ][ ``luminosity and density evolution\" model]{Aird10}, in excellent agreement with our evolution in the normalization of the Eddington ratio distribution presented above. We therefore conclude that our redshift evolution is the result of a shift in the accretion processes to lower characteristic Eddington ratios since $z\\sim1$.  This implies that any upper boundary to the Eddington ratio distribution would also move to lower \\Lamedd\\ as redshift decreases. Our results do not show evidence for such a redshift-dependent cutoff, although we probe only relatively low Eddington ratios, especially at low redshifts (see Figure \\ref{fig:leddfrac_zbin}), and do not include broad-line QSOs that may dominate the population at the highest Eddington ratios. Much larger samples are required to determine the presence and evolution of an upper boundary or turnover of our Eddington ratio distribution at high \\Lamedd. In fact, \\citet{Steinhardt10} have presented evidence for a ``sub-Eddington boundary\" in large samples of SDSS quasars at $0.2<z<4$, which depends on black hole mass and may evolve with redshift, although the connection to our results requires further investigation.  Our findings do not agree with the ``downsizing\" picture of AGN evolution, which attributes the luminosity evolution of the XLF to a decrease in the average mass of galaxies that host AGNs as time progresses, either due to a shut off of AGN activity in the highest mass galaxies or AGNs in lower mass galaxies ``turning on\" at later times \\citep[e.g.,][]{Barger05,Alonso-Herrero08,Hopkins09b}. This picture is not supported by our results: we find that the distribution of accretion rates is uniform for galaxies of all stellar masses at any given redshift. The entire distribution evolves with redshift, which we interpret as a shift to lower characteristic accretion rates with increasing cosmic time. Thus, at any given epoch, black hole growth is taking place in galaxies over the full range of stellar masses, but at later times the accretion rates are simply much lower, on average, and there is a relative lack of the most rapidly accreting sources. The evolution of the XLF is driven by this shift in accretion rate rather than the downsizing of black hole growth \\citep[see also][]{Shankar09}. However, as our results track the specific accretion rate relative to the host stellar mass (see Section \\ref{sec:specaccrate}), we are unable to constrain any underlying evolution of the distribution of black hole masses, which may evolve differently from the stellar masses of their host galaxies and could potentially exhibit a downsizing behavior.  While our results provide important constraints, it is currently unclear what physical mechanism could be responsible for fueling AGN activity and regulating the distribution and evolution of accretion rates. The observed drop in the global star formation rate density from $z\\sim1-2$ to $z\\sim0$ \\citep{Hopkins04} presumably reflects a drop in the availability of cold gas in galaxies to fuel star formation and potentially AGNs. However, gas must be driven to the central region of the host galaxy to fuel an AGN. The fueling of the AGN could be triggered by instabilities in the host galaxy caused either by secular evolution \\citep[e.g.,][]{Kormendy04}, galaxy-galaxy interactions, or major or minor mergers \\citep[e.g.,][]{Kauffmann00,Hopkins08}. Alternatively AGN activity could be fed by smooth cosmological cold gas accretion \\citep[e.g.,][]{Keres05}, stochastic accretion of gas within the galaxy \\citep[e.g.,][]{Hopkins06c}, or mass loss from evolved stars \\citep[e.g.,][]{Ciotti07,Kauffmann09}. Whichever process controls the fueling of AGNs must be relatively independent of the stellar mass of the host galaxy, at least over the range probed by our study. Significant additional study is required to fully reveal the processes that are driving AGN accretion activity.   \\subsection{The relationship between AGN activity and host galaxy color} \\label{sec:hostcolor}  We find that the prevalence of AGN activity is somewhat enhanced in galaxies with blue optical colors, once stellar-mass-dependent selection effects are fully accounted for. In addition, we show that the probability of hosting an AGN for galaxies throughout our stellar mass range is primarily determined by the Eddington ratio. This drives the observed increase in the fraction of galaxies hosting X-ray-\\emph{detected} AGNs as the stellar masses of the parent galaxy sample increase. Thus, galaxies throughout the color-magnitude diagram may host an AGN, but a higher fraction of AGNs are \\emph{detected} within the most luminous galaxies, which also tend to have redder colors. In fact, the probability of hosting an AGN above a given Eddington ratio is higher for galaxies within the blue cloud and the green valley.  Both \\citet{Xue10} and \\citet{Silverman09b} have previously investigated the fraction of AGNs as a function of host-galaxy color and stellar mass. Indeed, \\amend{\\citet{Silverman09b} highlighted the importance of stellar-mass-dependent selection effects and the potential bias in studies using optical flux-limited samples \\citep[e.g.,][]{Nandra07b,Coil09,Hickox09,Schawinski09,Georgakakis11}, which was later demonstrated explicitly by \\citet{Xue10}. The \\citet{Xue10} study also found that} moderate-luminosity X-ray AGNs are distributed over the full range of colors in stellar-mass-matched samples with no specific clustering in the color--magnitude diagram. However, their sample of AGNs at $0<z<1$ is comparatively small (83 X-ray-detected AGNs) and was combined over this entire redshift range where we have found strong evolution of the AGN fraction. In addition, they are mainly reliant on photometric redshift estimates, which have larger errors and a higher outlier fraction than our PRIMUS redshifts, \\amend{and probe to fainter optical magnitudes. These uncertainties will propagate to estimates of stellar masses and rest-frame colors, potentially blurring the color--magnitude diagram and masking} the relatively weak enhancement (factor $\\sim2$) of the AGN fraction that we see in galaxies in the blue cloud and green valley.  \\citet{Silverman09b} found that the X-ray AGN fraction at $0.5\\lesssim z \\lesssim1.0$ increases in galaxies with young stellar ages, traced by the $D_n(4000)$ spectral index, and in galaxies with blue rest-frame optical colors, consistent with our results. \\citet{Silverman09b} also presented evidence that the levels of ongoing star formation, traced by the  [O{\\sc ii}]$\\lambda$3727 emission-line luminosity, are higher in AGN host galaxies than in non-AGN galaxies with equivalent stellar masses.  \\citet[herafter S10]{Schawinski10} also studied the AGN fraction as a function of galaxy color and stellar mass. They used a local sample of galaxies from the SDSS ($z<0.05$), identified narrow-line AGNs within these galaxies using emission-line diagnostics and determined AGN luminosities (and thus Eddington ratios) from the intensity of the [O\\textsc{III}] emission line. S10 find that the fraction of detected AGNs increases with stellar mass for all galaxies, consistent with prior studies and our findings for X-ray-detected AGNs. An enhancement in the fraction of AGNs is seen in the green valley along with a significant lack of emission-line AGNs in galaxies on the red sequence. They also find a low fraction of AGNs in the blue cloud but note that for blue galaxies their AGN sample is incomplete for $L_\\mathrm{OIII}\\lesssim10^{41}$ \\ergs; at low luminosities the AGN emission lines are overwhelmed by the star formation precluding identification of the AGN. The blue galaxy sample is also generally at low stellar masses ($\\log \\mstel/\\msun\\lesssim10.5$). Using our best-fit results for $p(\\lamedd \\giv \\mstel,z)$ from Section \\ref{sec:eddratio} and assuming a bolometric correction of $L_\\mathrm{bol}\\approx450L_\\mathrm{OIII}$ \\citep{Kauffmann09}, we predict $\\lesssim1$\\% of galaxies with $9.5\\lesssim\\log\\mstel/\\msun\\lesssim10.5$ at $z<0.05$ to host AGN with $L_\\mathrm{OIII}>10^{41}$. Thus the results of S10 appear consistent with our findings, but the bias against the selection of AGNs in low-mass, star-forming galaxies in their work may have a significant effect.   \\subsection{The role of AGNs in galaxy evolution} \\label{sec:role}  What do our results imply about the connection between AGN activity and galaxy evolution? The fraction of galaxies hosting AGN rapidly decreases between $z\\approx1$ and $z\\approx0.2$ but follows the same distribution of Eddington ratios. We find strong evolution in the fraction of galaxies hosting AGN with redshift of the form $p(\\lamedd \\giv \\mstel,z) \\propto (1+z)^{3.5\\pm0.5}$. This evolution is very similar to the overall evolution in the star formation density, $\\dot\\rho_* \\propto (1+z)^{3.5\\pm0.2}$ for $0<z<1.2$ \\citep{Rujopakarn10}. This correspondence could be due to a relationship between the processes that fuel star formation and AGN: both could be fueled from the same gas supply,  or the AGN could be fueled indirectly via stellar mass loss from the stars within the galaxy \\citep[e.g.,][]{Norman88,Ciotti07}.   We also find a weak association between the prevalence of AGN activity and blue optical colors associated with current star formation activity. That said, AGNs with a wide range of Eddington ratios are found in galaxies of all colors. The \\emph{shape} of the distribution of Eddington ratios is roughly constant as a function of color. The enhancement of AGN activity in galaxies with blue colors could indicate that the distribution of Eddington ratios is shifted to higher \\Lamedd, on average. This could imply that red galaxies contain AGNs at later stages of their lifetimes when the rate of accretion has reduced and the luminosities have faded. It has been proposed that feedback from an AGN can quench star formation in galaxies, resulting in the red colors \\citep[e.g.,][]{Springel05,Hopkins06}. However, the most rapidly accreting SMBHs (highest Eddington ratio sources) in our sample are \\emph{not} exclusively in galaxies with current star formation (blue colors). Red colors and a lack of current star formation for a galaxy do not preclude substantial AGN accretion and significant SMBH growth. Equally, a high fraction of blue galaxies contain weakly accreting (low Eddington ratio) AGNs.  Thus, while AGN activity and star formation do appear to be correlated, we do not find evidence that AGNs play a direct role in shaping the global properties of their host galaxies or their evolution. However, galaxies and their central AGNs could co-evolve in that the AGN fraction, related to the duty cycle of AGNs, may be regulated by whatever process drives gas to the centers of galaxies and fuels black hole growth. As discussed in Section \\ref{sec:agntriggering} above, this process must result in a shift to lower characteristic accretion rates as cosmic time progresses. This does not prevent the triggering of AGN accretion in galaxies of any color but must lead to characteristically lower accretion rates in red galaxies \\emph{or} a reduction in the triggering rate for red galaxies. As with the redshift evolution of the Eddington ratio distribution, we cannot distinguish between a shift to characteristically higher \\Lamedd\\ for blue galaxy colors or an overall increase in the density of blue galaxies hosting AGNs at a given \\Lamedd.  We thus conclude that our results do not show evidence that feedback from moderate-luminosity AGNs plays a direct role in the evolution of galaxies, although star formation and AGN activity may have a common triggering mechanism. We note that our results have no bearing on any potential role that may be played by luminous QSOs in the evolution of galaxies.  \\subsection{Remaining issues and future prospects} \\label{sec:issues}  Despite our unprecedentedly large sample of galaxies and hard X-ray-selected AGNs with spectroscopic redshifts at $z\\sim0.2-1$, our results are dominated by Poisson uncertainties. This is due to the relatively small fraction of galaxies that host AGNs above a given X-ray luminosity ($\\sim 1$\\%$-10$\\% but highly dependent on stellar mass and redshift) as well as the depths of our X-ray data which impede detection of lower luminosity AGNs. We account for the Poissonian uncertainties in our analysis and are able to constrain functional forms that provide a good description of our data, although our discriminating power is limited by the source numbers. Indeed, the results of Section \\ref{sec:col} hint that the probability of hosting an AGN may have a complex color dependence that deviates from our global Eddington ratio dependence. However, due to our limited sample, we cannot split the data into many fine bins in redshift, stellar mass, X-ray luminosity and galaxy color at the same time to examine this behavior more closely.  We also note that our results are valid for the population of moderately obscured, moderate-luminosity AGNs probed by our data. This population appears to represent the majority of the AGN accretion density over our redshift range \\citep{Ueda03,Aird10}, and thus our results can shed light on the dominant processes that are responsible for triggering and regulating AGN activity over this epoch. Our results should not be extrapolated to the very low or high Eddington ratio regime where different accretion and fueling processes may be relevant. \\moreamend{Furthermore, the lowest stellar mass galaxies in our sample are only included at the lower end of our redshift range due to the redshift-dependent stellar mass cut, which is defined by the magnitude limits of our survey. This cut could introduce a bias in our results, although applying an additional cut of $\\log \\mstel/\\msun>10.5$ at all redshifts does not significantly alter our results (but does substantially increase the errors due to the smaller sample size). Conversely, higher stellar mass galaxies tend to be found at higher redshifts due to the larger volume surveyed, which could also introduce a bias. Both larger area surveys and deeper samples are required to address these issues.}  BLAGNs may dominate the population at the highest Eddington ratios, and thus their exclusion could also bias our derived Eddington ratio distribution. \\amend{To investigate the impact of BLAGNs, we apply a completeness correction term in our maximum-likelihood fitting based on the fraction of BLAGNs within the total sample of X-ray detected PRIMUS sources at $0.2<z<1.0$. We apply this correction as a function of X-ray luminosity (over the $42<\\log \\lx<44$ range) but assume no redshift evolution and implicitly assume that the host stellar masses of the BLAGNs have the same distribution as the non--broad-line sample. Using this basic correction, we find a slightly flatter slope, $\\gamma_E=0.62\\pm0.04$, which is consistent with the slope derived in Section \\ref{sec:eddratio} (within the errors) and indicates that the exclusion of BLAGNs does not introduce a significant bias.} \\amend{Nevertheless, this simple correction may not accurately account for this population.} BLAGNs may correspond to a specific phase in the lifetimes of AGNs, be powered by different accretion modes, or be a distinct population that is triggered and fueled by very different processes compared to our sample. \\amend{Indeed, BLAGNs dominate the X-ray detected population at high luminosities ($\\lx\\gtrsim10^{44}$ \\ergs), outside the range probed by our study and corresponding to the strong break seen in the XLF. Such sources may dominate the AGN population at the highest Eddington ratios and constraints on their accretion properties are vital to track the shape of the Eddington ratio distribution and any cutoff around the Eddington limit. Estimates of the SMBH mass for such sources are possible using high-resolution spectroscopy and scaling relations between the luminosity, the width of broad lines and the SMBH mass \\citep[e.g.,][]{Trump11}, which would allow estimates of the Eddington ratios and a more thorough investigation of the Eddington ratio distribution of BLAGNs. However, this is not possible with our PRIMUS spectra and is thus beyond the scope of this paper.}  Our hard X-ray selection may also bias our sample against the most absorbed AGNs ($\\NH\\gtrsim10^{23}$ cm$^{-2}$) and very heavily obscured, Compton-thick sources will be missed completely. While such sources may correspond to a substantial fraction of the AGN population \\citep[e.g.,][]{Gilli07,Draper09}, it is unclear if they represent a distinct phase of AGN growth and thus how their co-evolution with galaxies may differ to the moderately obscured population.  Our work shows that selection effects can severely bias samples of AGNs selected using different methods due to the interplay between the Eddington ratio, the observed AGN luminosity, and the stellar mass and star formation properties of the galaxy. However, by correcting for the various selection effects in our data, we are able to reveal the overall trends in the prevalence of AGNs and the dependence on redshift, stellar mass, X-ray luminosity or Eddington ratio. Thus, our results can strongly constrain models of AGN triggering and the co-evolution of AGNs and galaxies. Future work could push our studies to higher redshifts or compare our findings with relationships in the local universe. Alternatively, increased area coverage and X-ray depth could increase our sample size and allow a more thorough analysis of the relationship between luminosity, absorption properties, \\amend{bolometric corrections,} accretion rates, stellar masses, galaxy colors, star formation properties, and morphologies as a function of redshift over this key period of the universe, providing further insight into the physics of black hole growth and the relationship between black holes and their host galaxies. However, any such future studies must carefully account for stellar-mass-dependent selection effects and the X-ray sensitivity to shed light on the true distribution of AGN accretion.          In this paper we study how the probability of a galaxy hosting an AGN depends on stellar mass and color, as well as revealing the underlying distribution of accretion rates probed by the X-ray luminosities and Eddington ratios. We use an unprecedentedly large sample of galaxies with spectroscopic redshifts from PRIMUS, deep broad-band optical imaging, and X-ray data from \\chandra\\ or \\xmm.  Our main conclusions are as follows: \\begin{itemize}  \\item We find that for galaxies at $z=0.2-1$ the probability of hosting an AGN as a function of X-ray luminosity, $p(\\lx \\giv \\mstel,z)$, is a power-law function of \\Mstel\\ with slope $\\gamma_\\mathcal{M}$ that is independent of \\Lx. We also find that the distribution of X-ray luminosities at fixed stellar mass is a power law with slope $\\gamma_L$ that is independent of stellar mass. These results imply that while AGN activity is enhanced in higher stellar mass sources, the \\emph{distribution} of accretion rates is independent of mass.  \\item We find that the power-law slopes do not change with redshift, but there is strong evolution of the overall normalization. The probability of a galaxy hosting an AGN drops by a factor $\\sim 6$ between $z\\sim1$ and $z\\sim0.2$.  \\item We show that the probability of a galaxy hosting an AGN is defined by a universal Eddington ratio distribution, which is a power-law function of \\Lamedd\\ with slope $\\gamma_E=-0.65$ and is \\emph{independent} of the \\emph{stellar mass} of the host galaxy.  \\item Alternatively, these results can be interpreted in terms of \\emph{specific accretion rate}---the rate of black hole growth relative to the \\emph{stellar mass} of the host galaxy---rather than the true Eddington ratio. Thus, the distribution of \\emph{specific accretion rates} is described by a universal power-law function with slope $-0.65$ for all stellar masses.  \\item The Eddington ratio distribution evolves strongly with redshift with the form $(1+z)^{3.5\\pm0.5}$, similar to the rapid decrease in the total star formation rate density since $z\\sim1$. However, the drop in the overall black hole accretion rate density is driven by this decrease in the probability of hosting an AGN for galaxies of \\emph{all} stellar masses. We do not find a shift in the average stellar mass of a galaxy with an accreting black hole to lower masses or a shut down of AGN activity in the most massive galaxies as cosmic time progresses.  \\item We find that the probability of a galaxy hosting an AGN is mildly enhanced in galaxies with relatively \\emph{blue} colors by a factor $\\sim2$, once stellar-mass- and redshift-dependent selection effects are fully accounted for. However, AGN activity is taking place in galaxies of all colors. The shape of the distribution of AGN accretion rates does not appear to significantly vary between the blue cloud, green valley and red sequence, and the distributions span the full range of Eddington ratios.  \\item AGNs are \\emph{not} predominantly in red, massive galaxies. The previously observed increase in the AGN fraction with stellar mass is a selection effect that is driven by the \\Lamedd\\ distribution---a higher fraction of AGNs are detected in high stellar mass galaxies (with more massive black holes and generally redder colors) as sources accreting at low Eddington rates are intrinsically more luminous than in lower stellar mass galaxies. \\end{itemize}  Our results imply that the underlying physical processes that are responsible for triggering and fueling obscured AGN activity do not change between $z\\sim1$ and $z\\sim0.2$ but must decrease in frequency or shift to characteristically lower Eddington ratios. The same processes must take place in galaxies of all stellar masses and all colors, although activity appears to be enhanced in galaxies with current star formation. While AGN activity and star formation appear to be globally correlated, we do not find evidence that AGN feedback plays a direct and important role in the quenching and regulation of star formation in galaxies."
    },
    "1107/1107.1251_arXiv.txt": {
        "abstract": "We use the Excursion Set formalism to compute the properties of the halo mass distribution for a stochastic barrier model which encapsulates the main features of the ellipsoidal collapse of dark matter halos. Non-markovian corrections due to the sharp filtering of the linear density field in real space are computed with the path-integral technique introduced by Maggiore \\& Riotto \\cite{MaggioreRiotto2010a}. Here, we provide a detailed derivation of the results presented in \\cite{CorasanitiAchitouv2010} and extend the mass function analysis to higher redshift. We also derive an analytical expression for the linear halo bias. We find the analytically derived mass function to be in remarkable agreement with N-body simulation data from Tinker et al. \\cite{Tinker2008} with differences $\\lesssim 5\\%$ over the range of mass probed by the simulations. The excursion set solution from Monte Carlo generated random walks shows the same level of agreement, thus confirming the validity of the path-integral approach for the barrier model considered here. Similarly the analysis of the linear halo bias shows deviations no greater than $20\\%$. Overall these results indicate that the Excursion Set formalism in combination with a realistic modeling of the conditions of halo collapse can provide an accurate description of the halo mass distribution. ",
        "introduction": "Dark matter (DM) is an essential ingredient of the Standard Model of Cosmology. Observations provide strong evidence that about $90\\%$ of the matter in cosmic structures consists of such an invisible component \\cite{Spergel2003,Tegmark2004,Clowe2006,Massey2007}. Central to the DM paradigm is the idea that initial DM density fluctuations grow under gravitational instability fostering the collapse of baryonic matter. At late time, the gravitational infall becomes highly non-linear and DM particles violently relax into virialized objects, the halos. These are the building blocks in which cooling baryonic gas falls in to form the stars and galaxies that we observe today. Hence, the study of the halo mass distribution is of primary importance in Cosmology and an accurate modeling of the halo mass function has become essential to disclosing the complex mechanisms of the cosmic structure formation as well as probing the physics of the early universe. For example, recent studies have focused on the computation of the mass function for non-Gaussian initial conditions which are a direct signature of inflationary physics (see e.g. \\cite{Loverde2008,MaggioreRiotto2010c,Damico2011}).  The upcoming generation of galaxy cluster surveys (see e.g. \\cite{Pierre2011,ACT2010SPT2010}) will directly probe the distribution of massive halos, thus providing cosmological constraints complementary to those inferred from measurements of the Cosmic Microwave Background (CMB) radiation and cosmic distance indicators.  The properties of the DM halo mass distribution have been mainly investigated using high-resolution numerical N-body simulations. These studies have determined the halo mass function to a few percent accuracy level and provided insights on its redshift and cosmology dependence \\cite{Tinker2008,Courtin2010,Crocce2010,Suman2011}. In contrast, the development of purely theoretical studies has lagged behind. To date, a robust theoretical description of the halo mass function is still missing. For instance, we do not have a complete understanding of the relation between the conditions that lead to the formation of halos and the form of the mass function as obtained from N-body simulations. The numerical analysis usually limits to providing fitting formulae which depends on several ad-hoc parameters.  The seminal work by Press \\& Schechter (PS) \\cite{PressSchechter1974} is the first to attempt a derivation of the halo mass function from the statistical properties of the initial dark matter density fluctuation field. The idea is that halos form in regions in which the linearly extrapolated density field, smoothed on a given scale, lies above a given critical threshold of collapse, such as that predicted from the spherical collapse model \\cite{GunnGott1972}. The PS formula explicitly depends on such a threshold, which introduces an exponential cut-off in the high-mass end as confirmed by N-body simulation analysis. However, the PS computation suffers of an inconsistent behavior when the variance becomes arbitrarily large. In this asymptotic regime it is natural to expect that the mass fraction tends to $1$ (i.e. the higher the amplitude of the density fluctuations the larger is the fraction of mass in halos). However, in the PS case one finds $1/2$, thus suggesting that half of the mass in halos has been somehow miscounted. This is the so called `cloud-in-cloud' problem (see e.g. \\cite{PeacockHeavens1990} and discussion therein).  The formulation of the Excursion Set theory by Bond et al. \\cite{Bond1991} has provided the Press-Schecther approach with a powerful mathematical formalism in which the computation of the mass function is reduced to solving a stochastic calculus problem (for an exhaustive review of the formalism see \\cite{Zentner2007}). As shown in \\cite{Bond1991} the smoothed density fluctuation field behaves as a stochastic variable performing a random walk as function of the smoothing scale. Then, halos are associated to random trajectories which first cross a critical density threshold of collapse. It is the requirement of first-crossing which provides the solution to the `cloud-in-cloud' problem. Nevertheless, the computation of the halo mass function for realistic halo mass definitions has remained a challenging task. In fact, in the Excursion Set theory the mass of a halo depends on the form of the filter function, with the latter determining the behavior of the random walks. For realistic mass definitions the associated filters cause the random walks to be correlated, consequently the mass function can be estimated only through numerical Monte Carlo simulations \\cite{Bond1991,Percival2001}. Because of this, a thorough systematic comparison of the Excursion Set mass function obtained accounting for such correlations against N-body simulation data has never being performed.  A major breakthrough in this direction has been recently made by Maggiore \\& Riotto \\cite{MaggioreRiotto2010a} who have introduced path-integral techniques to perform an analytic computation of the halo mass function for generic filters. This allows us to consistently compare the Excursion Set model predictions with N-body data as well with astrophysical measurements of halo abundances.  Here, we compute the halo mass function and the linear halo bias for a barrier model which captures the main features of the ellipsoidal collapse of halos over a large range of masses. In this paper we also provide a detailed derivation of the results presented in \\cite{CorasanitiAchitouv2010}.  The paper is organized as follows. In Section~\\ref{excform} we briefly introduce the Excursion Set formalism, in Section~\\ref{barrmod} we discuss the modeling of the non-spherical collapse of halos and the computation of the mass function. In Section~\\ref{nonmark} we present the calculation of the corrections due to the filter function. In Section~\\ref{numan} we discuss the results of the comparison with N-body simulation data. In Section~\\ref{secbias} we present the computation of the halo bias. Finally, we present our conclusion in Section~\\ref{conclu}. ",
        "conclusions": "The Excursion Set formalism provides a powerful mathematical framework which allows us to perform a theoretical computation of the halo mass function from a limited set of initial assumptions. These must involve the statistics of the linear density fluctuation field as well as a stochastic barrier model of the halo collapse conditions. In addition, such calculation needs to be implemented with a path-integral evaluation of the corrections due to the filtering of the linear density field associated with a realistic mass definition. Such an approach allows for a consistent model comparison with N-body simulation data.  Here, we have derived an analytical expression for the mass function and linear bias in the case of a diffusive barrier model with linearly drifting average. This model well approximates the main features of the ellipsoidal collapse. We have found a remarkable agreement with N-body simulation data with differences $\\approx 5 \\%$ over a large range of masses. Such an agreement is due to the competing effects of the barrier average drift at small masses and of the diffusion in the high-mass end. This has important phenomenological implications especially in the study of primordial non-Gaussianity. In fact, several studies have estimated the halo mass function in the case of non-Gaussian initial conditions assuming the spherical collapse models (see e.g. \\cite{MatarreseVerdeJimenez2000,Loverde2008,MaggioreRiotto2010c}). However the comparison with non-Gaussian N-body simulations has shown large deviations in the low mass range compared to large masses (see e.g. Fig. 1 in \\cite{GiannantonioPorciani2010}). In the light of our results, it is plausible that such discrepancies may be attributed to the non-spherical collapse of halos. The inclusion of a simple diffusive barrier model with linearly drifting average for a non-Gaussian linear density field could resolve or alleviate the problem.  Our results suggest a number of directions which warrent further investigation. Firstly, it will be insightful to derive the statistical properties of the fuzzy barrier for a given ellipsoidal model as function of the variance of the linear density field. This will provide theoretical predictions for $\\beta$ and $D_B$ which can be confronted with numerically calibrated values for different redshifts and cosmologies, and it will allow us to accurately testing the modeling of the halo collapse condtions. On the other hand, in the upcoming years several observational campaigns will probe the halo abundance through galaxy cluster surveys. The mass function derived here can be used to perform a data analysis of the barrier model parameters, thus providing information on the collapse of DM halos which has been previously unforseen."
    },
    "1107/1107.3943_arXiv.txt": {
        "abstract": "We perform a Bayesian parameter inference in the context of resonantly damped transverse coronal loop oscillations. The forward problem is solved in terms of parametric results for kink waves in one-dimensional flux tubes in the thin tube and thin boundary approximations. For the inverse problem, we adopt a Bayesian approach to infer the most probable values of the relevant parameters, for given observed periods and damping times, and to extract their confidence levels. The posterior probability distribution functions are obtained by means of Markov Chain Monte Carlo simulations, incorporating observed uncertainties in a consistent manner. We find well localized solutions in the posterior probability distribution functions for two of the three parameters of interest, namely the Alfv\\'en travel time and the transverse inhomogeneity length-scale. The obtained estimates for the Alfv\\'en travel time are consistent with previous inversion results, but the method enables us to additionally constrain the transverse inhomogeneity length-scale and to estimate real error bars for each parameter. When observational estimates for the density contrast are used, the method enables us to fully constrain the three parameters of interest. These results can serve to improve our current estimates of unknown physical parameters in coronal loops and to test the assumed theoretical model. ",
        "introduction": "Magnetohydrodynamic (MHD) seismology stands as one of the few indirect methods for the determination of difficult to measure physical parameters in solar coronal structures. It relies on the combined use of observed oscillatory properties of MHD waves in solar atmospheric magnetic and plasma structures together with theoretically obtained wave properties. It was first suggested by  \\citet{uchida70} and \\cite{REB84}, in the coronal context, and by \\citet{TH95}, in the context of solar prominences. The last years increase in the number and quality of observations of wave activity in the solar atmosphere and the refinement of  theoretical models have allowed the practical implementation of this technique.  In the context of coronal loops, coronal seismology has allowed the estimation and/or restriction of  relevant parameters such as the magnetic field strength \\citep{Nakariakov01}, the Alfv\\'en speed \\citep{temury03,Arregui07,GABW08}, the transversal density structuring \\citep{verwichte06}, or the coronal density scale height \\citep{AAG05,Verth08}.  Of particular relevance has been the use of the concept of period ratios as a seismological tool, first pointed out by  \\cite{AAG05,GAA06}, and reviewed by \\cite{andries09}.   Most of the efforts in this area have been concentrated on the phenomenon of quickly damped transverse oscillations in coronal loops, first reported by \\citet{Aschwanden99} and \\cite{Nakariakov99}. These oscillations are interpreted in terms of linear MHD kink waves of a magnetic flux tube, a wave mode with mixed fast and Alfv\\'en character, its Alfv\\'enic nature being dominant in and around the resonant position \\citep{Goossens09}, where the global eigenmode frequency matches the local Alfv\\'en frequency. Starting with the simplest model that considers the fundamental transverse oscillation of a magnetic flux tube \\citep{ER83} several model improvements have included other effects, such as the curvature of coronal loops, density stratification, or the departure from circular cross section of the tubes \\citep[see][for a recent review]{Ruderman09}. All these new ingredients have been seen to produce second order effects on the main wave properties.   The advancement in seismology of kink modes in coronal loops is reviewed by \\cite{Goossens08}. \\cite{Nakariakov01} used the observed periods and theoretical estimates of the periods, based on the long wavelength  approximation  for a uniform coronal loop model, to derive estimates for the magnetic field strength, by making assumptions on the density. \\cite{GAA02} used the observed damping rates and theoretical values of the damping rates, based on the thin boundary approximation, to derive estimates for the radial inhomogeneity length-scale, once the density contrast was assumed. \\cite{Aschwanden03b} used the observed damping rates and the damping rates computed by \\cite{tom04b}, outside the thin boundary approximation, to compare the theoretically predicted density contrasts to estimates obtained indirectly taking into account effects such as the subtraction of the background flux, the line-of-sight integration of the emission measure, the spatial smearing due to the transverse motion, and the point-spread function of the instrument. In their study, the internal density could only be determined as a function of the external density.   The main limitation of these studies is that the determination of the magnetic field strength \\citep{Nakariakov01} or transverse inhomogeneity length-scale \\citep{GAA02} is only possible if a given value for the density contrast is adopted. Otherwise, an infinite number of solutions to the inverse problem arise. This was shown in the numerical and analytic seismological inversions by \\cite{Arregui07} and \\cite{GABW08},  that combine the observational information on both periods and damping times for resonantly damped kink modes in a consistent manner.  These inversions enable to infer information about  both the internal Alfv\\'en speed and the transverse density structuring.  Their approach is to make no assumption on the particular value of any of  the physical parameters of interest. Their ranges of variation, compatible with observations, are instead obtained \\citep{GABW08}. Following this approach, a complete solution to the inverse problem is obtained, from which general properties and liming cases can be studied. The solution gives rise to an infinite number of equally valid equilibrium models that explain observations, characterized by three parameters: the density contrast, the transverse inhomogeneity length-scale, and the Alfv\\'en speed. In addition, in the inversion techniques presented by \\citet{Arregui07} and \\citet{GABW08}, it is not straightforward to devise a consistent method to compute uncertainties on the inferred parameters from the measurement errors on observed periods and damping rates.  Bayesian inversion techniques for parameter inference can help us to overcome these limitations and are applied in this paper to the determination of physical parameters in oscillating coronal loops. These techniques have proven to be useful in a number of areas of physical sciences \\citep{Gregory05}, including solar physics. In the context of inversion of Stokes profiles, \\cite{andres07} apply Bayesian techniques to analyze the performance of a given radiative transfer model to fit a given observed Stokes vector aiming at obtaining information about the thermodynamic and magnetic properties of solar and stellar atmospheres \\citep[see also][]{andres11}. \\cite{Marsh08} use a Bayesian probability-based approach to the problem of detecting and parameterizing oscillations in the upper solar atmosphere, in order to determine the number of oscillations present, and their properties.  In the context of coronal heating, \\citet{adamakis10} employ Bayesian model comparison techniques  to determine the preferred heating location along coronal loops. Finally, \\cite{Ireland10} have recently presented an automated detection technique for oscillating regions in the solar atmosphere based on  Bayesian spectral analysis of time series and image filtering.  In this paper, a Bayesian MHD seismology inversion technique is presented and the results from its application to coronal loops are described. The layout of the paper is as follows. Section~\\ref{forinv} describes the physical model and the forward and inverse problems.  In Section~\\ref{tech}, the details of the developed Bayesian technique are given. The analysis and results of its application to synthetic data and to real coronal loop oscillations are discussed in Section~\\ref{results}. Finally, in Section~\\ref{conclusions}, our conclusions are presented. ",
        "conclusions": "We have applied a Bayesian parameter inversion technique to the determination of unknown physical parameters in transversely oscillating coronal loops.  The model considers  coronal loops as one-dimensional cylindrically symmetric density enhancements with a radial density variation. The forward problem reduces to an algebraic set of equations for two observables as a function of three parameters. A Markov Chain Monte Carlo algorithm is used to sample the posterior probability distribution for the density contrast, the transverse inhomogeneity length-scale, and the internal Alfv\\'en travel time. We find that the latter two parameters can be very well determined, while the density contrast cannot in general be constrained.  Density contrast estimations are challenging in both observations and MHD seismology inversions.  By using synthetic data to perform the inversion under controlled conditions, different prior information for the density contrast has been tested. A uniform prior distribution is unable to correctly infer the parameters of interest, unless some a priori information on density contrast is known, since it produces biased results  when extended ranges are considered. On the other hand, a Jeffreys type prior is more appropriate to perform the inversion when information on densities is lacking or uncertain.  If additional information on density contrast is available from observations, the three unknowns can be very well determined, by using a Gaussian prior centered on the observed density contrast value. The reliability of the inversion then depends on the reliability of the density contrast  measurement.   The advantage of the proposed technique with respect to \\cite{Nakariakov01,GAA02, Aschwanden03b} is that it uses the information contained in both the period and the damping rate of oscillating coronal loops. In addition it makes no assumption on the particular value of the unknown parameters. The advantage with respect to the analytic and numerical inversions by \\cite{Arregui07,GABW08} is that, first,  it enables us to  constrain  the transverse inhomogeneity length-scale and the density contrast. Second, the method incorporates  confidence levels and error bars consistently calculated  from the uncertainties on observed wave properties. Bayesian inference gives  a measure of degree of belief  about a proposition (set of parameters)  given some conditionals (observed data), once a probability model is set up. In that sense it is appropriate to perform inversions by making the fewer possible assumptions in the comparison to observations, as proposed by \\cite{Arregui07} and \\cite{GABW08}.  Our study is based on a theoretical model that assumes the classic straight magnetic flux tube representation for coronal loops and resonant absorption as damping mechanism. If any of these assumptions deviates from reality, then this will introduce errors into the inversion presented in this paper. For instance, \\cite{demoortelpascoe09} have shown that the magnetic field strength inside a three-dimensional numerical coronal loop model may substantially differ from that inferred using classic seismology inversions in straight magnetic flux tubes.  The presented inversion technique can be applied to the determination of  physical parameters in other magnetic and plasma structures of the solar atmosphere that display oscillatory dynamics.  One example are transversely oscillating prominence threads, provided a large part of the magnetic flux tube is filled with dense plasma \\citep[see][]{Soler102dthread}. In this case, and because of the large density contrast typical of prominence plasmas, with internal densities two orders of magnitude larger than coronal densities, this parameter becomes irrelevant  to the inversion process \\citep{AB10}. If we apply our algorithm to a case with $P=3$ min and $\\tau_d=9$ min (with 10\\% error), by considering large density contrast limits in the prior information for this parameter, we  find estimates of $l/R=$0.22$^{+0.03}_{-0.03}$, $\\tau_{\\rm Ai}=$129$^{+12.8}_{-13.1}$, in excellent agreement with the analytic inversion approximation of  $l/R=0.21$, $\\tau_{\\rm Ai}=126.9$.  Although it may seem that in this case, as in the coronal loop inversion with Gaussian prior information, a two parameter-two observable problem seems to be fully constrained, we should be aware that we are dealing with incomplete information, because of uncertainties from measured quantities. The advantage of our method then lies in the consistent propagation of errors.    One of the most appealing applications of Bayes' rule is model comparison, where the evidence in Equation~(\\ref{bayes}) plays the fundamental role. Provided different damping mechanisms could be properly parameterized with similar physical model parameters and observables, a Bayesian model comparison could be performed in order to shed light on the still highly debated mechanism that produces  the damping of transverse oscillations in coronal structures."
    },
    "1107/1107.1946_arXiv.txt": {
        "abstract": "Motivated by recent observational results that focus on high redshift black holes,  we explore the effect of scatter and observational biases on the ability to recover the intrinsic properties of the black hole population at high redshift.  We find that scatter and selection biases can hide the intrinsic correlations between black holes and their hosts, with `observable' subsamples of the whole population suggesting, on average, positive evolution  even when the underlying population is characterized by no- or negative evolution.  We create theoretical mass functions of black holes convolving the mass function of dark matter halos with standard relationships linking black holes with their hosts.  Under these assumptions, we find that the local $M_{\\rm BH}-\\sigma$ correlation is unable to fit the $z=6$ black hole mass function proposed by Willott et al. (2010), overestimating the number density of all but the most massive black holes. Positive evolution or including scatter in the $M_{\\rm BH}-\\sigma$ correlation makes the discrepancy worse, as it further increases the number density of observable black holes.  We notice that if the $M_{\\rm BH}-\\sigma$ correlation at $z=6$ is steeper than today, then the mass function becomes shallower. This helps reproducing the mass function of $z=6$ black holes proposed by Willott et al. (2010). Alternatively, it is possible that very few halos  (of order 1/1000) host an active massive black hole at $z=6$, or that most AGN are obscured, hindering their detection in optical surveys. Current measurements of the high redshift black hole mass function might be underestimating the density of low mass black holes if the active fraction or luminosity are a function of host or black hole mass.  Finally, we discuss physical scenarios that can possibly lead to a steeper  $M_{\\rm BH}-\\sigma$ relation at high redshift. ",
        "introduction": "The constraints on black hole masses at the highest redshifts currently probed, $z\\simeq6$, are few, and seem to provide conflicting results.  (i) There seems to be little or no correlation between black hole mass and velocity dispersion, $\\sigma$ \\citep{Wang2010} in the brightest radio-selected quasars, (ii) typically black holes are `over-massive' at fixed galaxy mass/velocity dispersion compared to their $z=0$ counterparts \\citep[e.g., Walter et al. 2004; at lower redshift see also][]{Mclure2004,Shields2006,Peng2006b, Decarli2010, Merloni2010, Woo2008}, but (iii) analysis of the black hole mass/luminosity function  and clustering suggests that either many massive galaxies do not have black holes, or these black holes are less massive than expected \\citep[W10 hereafter]{Willott2010}.  As a result of point (ii), most authors propose that there is {\\it positive} evolution in the $M_{\\rm BH}-$galaxy relationships, and quantify it as a change in {\\it normalization}, in the sense that at fixed galaxy properties (e.g. velocity dispersion, stellar mass), black holes at high redshift are more massive than today. For instance, Merloni et al. (2010) propose that $M_{\\rm BH}-M_*$ evolves with redshift as $(1+z)^{0.68}$ while \\cite{Decarli2010} suggest $(1+z)^{0.28}$.   Point (iii) above, however, is inconsistent with this suggestion unless only about $1/100$ of galaxies with stellar mass $\\simeq 10^{10}-10^{11} $ \\msun ~at $z=6$ host a black hole (W10). These galaxies are nonetheless presumed to be the progenitors of today's massive ellipticals, which typically host central massive black holes.  When inferences on the population of massive black holes at the highest redshift are made, we have to take into consideration two important selection effects (see Lauer et al. 2007b). First, only the most massive black holes, powering the most luminous quasars, can be picked up at such high redshifts \\citep{Shen2008, Vestergaard2008}. Second, as a result of the limited survey area of current imaging campaigns, only black holes  that reside in relatively common galaxies can be recovered. Taken together, these biases imply  that the observable population of black holes at high redshift will  span a narrow range of masses and host properties (see also Adelberger et al. 2005, Fine et al. 2006).  In this paper, we explore the impact of these observational biases on attempts to recover the intrinsic properties of the black hole population. Our calculations are based on simple models grounded on empirical relations measured at much lower redshift, and therefore our results should be treated with caution. The aim of this paper is only to highlight  the effects of the different factors that can influence the measurement of the intrinsic properties of the black hole population at high redshift.  In section 2 we describe how we generate Monte Carlo realizations of the $M_{\\rm  BH}-\\sigma$ relation at $z=6$  varying the slope and normalization. We then select `observable' systems from these samples, considering both `shallow' or `pencil beam' surveys, and test how well we can recover  the parameters  of the $M_{\\rm  BH}-\\sigma$ relation from the `observable' systems.   In section 3 we discuss theoretical mass functions of black holes derived from the mass function of dark matter halos and various assumptions for the $M_{\\rm BH}-\\sigma$ relationship. Using these results, we test what assumptions can reproduce the black hole mass function derived by W10.  We also discuss (section 4) why obtaining constraints on the average accretion rates and active fraction of black holes as a function of host mass is crucial to our understanding of the high-redshift massive black hole population. Finally, in section 5 we propose a simple theoretical framework that leads to selective accretion onto black holes in a manner that reconciles the observational results (i)-(ii) and (iii) above. ",
        "conclusions": "Recent observational results that focus on high-redshift black holes provide seemingly conflicting results. In particular: \\begin{itemize} \\item[(i)] there seems to be little or no correlation with velocity dispersion, $\\sigma$ in the brightest radio-selected quasars, \\item[(ii)] typically black holes are `over-massive' at fixed galaxy mass/velocity dispersion compared to their $z=0$ counterparts, \\item[(iii)] clustering and analysis of the mass/luminosity function suggest that either many massive galaxies do not have black holes, or these black holes are less massive than expected. \\end{itemize}  To try and understand these observational results, we explore the role of scatter and observational biases in recovering the intrinsic properties of the black hole population.  We generate Monte Carlo realizations of the $M_{\\rm  BH}-\\sigma$ relation at $z=6$, varying the slope and normalization, and select `observable' systems from these samples, considering either `shallow' or `pencil beam' surveys. We test how well we can recover the parameters  of the $M_{\\rm  BH}-\\sigma$ relation from the `observable' systems only.  We then create theoretical mass functions of black holes from the mass function of halos and $M_{\\rm BH}-\\sigma$ and test what assumptions can reproduce the mass function derived by W10. Our techniques are very simplified and we use empirical correlations that are not guaranteed to hold at all masses and reshifts.  Therefore one should not interpret our results as solutions to the three conflicting points mentioned above, but rather regard them as a way for understanding how different physical parameters may affect black hole related quantities and their measurements.  Our results can be summarized as follows:  \\begin{itemize} \\item Scatter and bias selections can hide the intrinsic correlations between holes and hosts. When selecting within a small range in black hole and galaxy masses, at the high-mass end, the scatter washes out correlations (see point (i) above), and most of the selected systems will tend to lie above the underlying correlation. The correlations recovered from `observable' sub-samples of the whole population can therefore suggest positive evolution  even when the underlying population is characterized by no or negative evolution. \\item The slope and normalization of the local $M_{\\rm BH}-\\sigma$ correlation are unable to produce a black hole mass function compatible with W10, as the theoretical mass function greatly overstimates the density of black holes with $M_{\\rm BH}<10^8$. The discrepancy can be minimized if very few halos (of order 1/1000) host a massive black hole or an AGN at $z=6$, or most AGN at these redshift are obscured. \\msun. \\item If the $M_{\\rm BH}-\\sigma$ correlation were steeper at $z=6$, then at fixed black hole mass high mass black holes would reside in comparatively less massive galaxies than in the $\\alpha=4$ case. Their number density is therefore increased. Viceversa, low mass black holes would be hosted in comparatively larger galaxies (compare red and yellow lines in Figure~1) with a lower space density. This effect helps reproducing the mass function of $z=6$ black holes proposed by W10. \\item On the other hand, scatter in the $M_{\\rm BH}-\\sigma$, at the level of what is observed locally, exacerbates the discrepancy, as it increases the number density of black holes at $M_{\\rm BH}>10^7$ \\msun. Any type of positive evolution of the $M_{\\rm BH}-\\sigma$ exacerbates this discrepancy. \\item Analysis of AGN samples might be underestimating the black holes mass function at low masses if the active fraction or luminosity are a function of host or black hole mass. \\end{itemize}  In the near future the synergy of {\\it JWST} and {\\it ALMA} can zoom in on quasars and their hosts respectively informing us of their relationship and how the $M_{\\rm BH}-\\sigma$ relation is established, or how the accretion properties depend on the black hole or halo mass.  In the near-IR, {\\it JWST} will have the technical capabilities to detect quasars at $z\\simgt 6$ down to a mass limit as low as $10^5-10^6$ \\msun, owing to its large field of view and high sensitivity. At the expected sensitivity of {\\it JWST}, $\\simeq 1$ nJy, almost $7\\times 10^3$ deg$^{-2}$ sources at $z> 6$ should be detected \\citep{Salvaterra07}.  At the same time, the exquisite angular resolution and sensitivity of {\\it ALMA} can be used in order to explore black hole growth up to high redshift even in galaxies with high obscuration and active star formation. To date the best studies of the hosts of $z\\simeq 6$ quasar have been performed at cm-wavelength \\citep{Walteretal2004,Wang2010}. The best studied case is J1148+5251 at z = 6.42.  The host has been detected in thermal dust, non-thermal radio continuum, and CO line emission \\citep{Bertoldi2003,Carilli2004,Walteretal2004}.  {\\it ALMA} will be able to detect the thermal emission from a source like J1148+5251 in a few seconds at sub-kpc resolution \\citep{Carilli2008}. On a similar time-frame Dark-Energy oriented survey will provide an enormous amount of quasar data as ancillary science (e.g. DES, LSST). Coupling the information we derive from these extremely large yet shallow surveys, with that derived from deep pencil beam surveys will undoubtedly deepen our understanding of the growth of high-redshift black holes.  In the meantime, we need to develop dedicated cosmological simulations of black hole formation and early growth that can aid the interpretation of these data.  The suggestion that the accretion rate of massive black holes depends on their environment, (through the host halo and its cosmic bias) must be tested with cosmological simulations that implement physically-motivated accretion and feedback prescriptions. We also need to derive predictions for the occupation fraction of black holes in galaxies based on black hole formation models (Bellovary et al. in prep.).  This will be a significant improvement over current simulations of black hole cosmic evolution that typically place black holes in halos growing above some threshold mass, typically $\\sim 10^{10}$ \\msun, leading to a trivial occupation fraction function.  There is no strong physical reason to believe that all and only galaxies with mass $>10^{10}$\\msun~ host massive black holes in their centers.  In this paper we focused on the very high-redshift Universe, $z\\simeq 6$. Although this redshift range is not a special place, the concurrence of theoretical arguments and observational constraints allow us to make simplifying assumptions that are not expected to be valid at later times. For instance,  a timescale argument requires black holes to grow fast to reach the masses probed by current luminosity functions. In turn, this argument, coupled to current observational constraints suggest that the most luminous quasars accrete close to Eddington, and that both active fraction and occupation fractions must be of order unity at the high-luminosity, high-mass end. This is not true at $z=2$, where more variables enter into play, and make the analysis less constraining.  An example of a detailed study that connects the mass function of black holes derived from the Press \\& Schechter formalism to that derived from the luminosity function is presented in Croton (2009)."
    },
    "1107/1107.0201_arXiv.txt": {
        "abstract": "We report on the first broad-band (1--200\\,keV) simultaneous \\chandra-\\integral observations of the recently discovered hard X-ray transient \\mysou~that took place  on 2011, March 22, about two weeks after the source discovery. The source had an average absorbed 1--200\\,keV flux  of about 8$\\mathrm{\\times 10^\\mathrm{-10}~erg~cm^\\mathrm{-2}~s^\\mathrm{-1}}$. We  extracted a precise X-ray position of \\mysou, $\\alpha_\\mathrm{J2000}$=17$^\\mathrm{h}$ 17$^\\mathrm{m}$ 42$^\\mathrm{s}$.62, \\mbox{$\\delta_\\mathrm{J2000}$= $-$36$^{\\circ}$ 56$^{\\prime}$ 04.5$^{\\prime \\prime}$} (90\\% uncertainty of 0.6$^{\\prime\\prime}$). We also report \\swift, near infrared and quasi simultaneous radio follow-up observations. With the multi-wavelength information at hand, we propose \\mysou~is a low-mass X-ray binary, seen at high inclination, probably hosting a black hole. ",
        "introduction": "\\label{sec:intro} \\setcounter{footnote}{0}   On 2011 March 15 (MJD 55635), \\integral discovered the new hard X-ray transient IGR~J17177$-$3656 \\citep{frankowski11}. The IBIS/ISGRI  spectrum (20--200\\,keV) could be well described by a power-law with photon index 1.8$\\pm$0.3 and a flux of about 20\\,mCrab. The source was also marginally detected  in JEM-X  in the 10--20\\,keV band (about 8\\,mCrab) but was not detected in the 3--10\\,keV band with a 3$\\sigma$ upper limit of 5\\,mCrab.  Following the \\integral discovery, a \\swift ToO was performed on 2011 March 16 \\citep{zhang11}. A refined \\swift position was reported (90\\% uncertainty of  2.1$^{\\prime \\prime}$), consistent with the \\integral one. A combined, though non simultaneous, \\swift/XRT and \\integral/IBIS/ISGRI spectral fit with an absorbed power-law model yielded $N_\\mathrm{H}$=\\errtwo{3.9}{0.5}{0.4}$ \\times 10^\\mathrm{22}~\\mathrm{cm^\\mathrm{-2}}$  and photon index $\\Gamma$=1.5$\\pm$0.2 indicating a hard state with additional absorbing column density intrinsic to the source \\citep{zhang11}. Based on the \\swift position and uncertainty, a search through the VizieR database resulted only in one match: 2MASS~J17174269$-$3656039 (K=12.9) located at about 1.3$^{\\prime \\prime}$ from the \\swift position of \\cite{zhang11}.\\\\  Our \\chandra ToO was performed on 2011 March 22, and thanks to the excellent \\chandra location accuracy, an X-ray  source position  with a 0.6$^{\\prime \\prime}$ (90\\%) uncertainty was reported \\citep{paizis11}. The new \\chandra-based position, not consistent with the \\swift  one ($\\sim$4.8$^{\\prime \\prime}$ away, but see Section~\\ref{sec:swiftresults}), probably  ruled out the 2MASS association (1.05$^{\\prime \\prime}$ away) and  allowed the scientific community as well as members of our team to search for NIR and radio counterparts \\citep[this paper;][]{corbel11,torres11,rojas11}. \\\\ The nature of \\mysou~is still to be unveiled and in this paper we describe our multi-wavelength campaign and contribution to its investigation. ",
        "conclusions": "\\label{sec:discussion} In order to investigate the  nature of \\mysou, we note that \\mysou~is located in the Galactic plane at (\\emph{l},\\emph{b})=(350$^{\\circ}$.1, 0$^{\\circ}$.51). Its X-ray emission together with its quasi-simultaneous radio emission, suggest that we are dealing with an X-ray binary, rather than with an AGN seen through the Galactic plane. From our least absorbed spectrum in Table~1, we obtain for the \\emph{un-absorbed} fluxes, relative to \\chandra normalization, F$_{\\mathrm{3-9\\,keV}}$=7.5$\\times \\mathrm{10^{-11}~erg~cm^{-2}~s^{-1}}$ and F$_{\\mathrm{3-200\\,keV}}$=1$\\times \\mathrm{10^{-9}~erg~cm^{-2}~s^{-1}}$. Using the relation found by \\cite{merloni03} that defines a \"fundamental plane\" in the three-dimensional (L$_{\\mathrm{< 10\\,keV}}$, L$_{\\mathrm{R}}$, M) space, and our quasi-simultaneous X-ray (un-absorbed 3--9\\,keV) and  radio emission, we see that we would need to place the source at about 100\\,Mpc to obtain a 3--9\\,keV luminosity of $\\sim$$\\mathrm{10^{44}~erg~s^{-1}}$, obtaining an estimated mass of about 4.4$\\times$$10^\\mathrm{4}$ $\\msun$ which is too low for an AGN. Placed at 1\\,Gpc, the source would have a luminosity as high as L$_\\mathrm{X}$$\\sim$$\\mathrm{10^{46}~erg~s^{-1}}$, still not reaching $10^\\mathrm{6}$ $\\msun$ (4.7$\\times$$10^\\mathrm{5}$ $\\msun$), which is extremely unlikely for an AGN scenario. \\\\ On the other hand, considering a binary at a distance of 8\\,kpc would result in an un-absorbed luminosity  L$_{\\mathrm{3-9\\,keV}}$=5$\\mathrm{\\times 10^{35}~erg~s^{-1}}$ and L$_{\\mathrm{3-200\\,keV}}$=7$\\mathrm{\\times 10^{36}~erg~s^{-1}}$ which is compatible with an X-ray binary luminosity. \\\\ Although the \\chandra observation occurred already in the declining phase of the outburst, its flux is comparable to the peak seen by \\swift (see Fig.\\ref{fig:lcr_all}), hence \\mysou~likely belongs to the so-called class of Very Faint X-ray Transients (VFXT), i.e., transients showing outbursts with low peak luminosity 10$^\\mathrm{34}$--10$^\\mathrm{36}$ erg~s$^\\mathrm{-1}$ in 2--10\\,keV \\citep{wijnands06}.\\footnote{Indeed even considering the \\swift peak flux according to \\cite{zhang11} (occurring on MJD 55636.8), we obtain a non absorbed F$_{\\mathrm{2-10\\,keV}}$=8.2$\\mathrm{\\times 10^{-11}~erg~cm^{-2}~s^{-1}}$, leading to L$_{\\mathrm{2-10\\,keV}}$=6$\\mathrm{\\times 10^{35}~erg~s^{-1}}$ at 8\\,kpc.}\\\\ VFXT are believed to be the faintest known accretors, and are very likely a non-homogeneous class of sources. It is likely that most of them are neutron stars and black holes  accreting matter from a low mass companion \\citep{wijnands06} and it has been found that a significant fraction of them ($\\sim$1/3) have exhibited type-I X-ray bursts \\citep[][and references therein]{delsanto2010}. This is currently not the case for \\mysou, and the black hole possibility, though more difficult to infer, is still open (see Section \\ref{sec:radio}).    \\subsection{LMXB or HMXB?}\\label{sec:HorL}   In order to speculate over the nature of the companion (hence high-mass X-ray binary, HMXB, versus Low-mass X-ray binary, LMXB), we have to consider the $K_\\mathrm{s}$ limits obtained in quiescence, so that we are not contaminated by the disc X-ray to NIR re-processing, as it would be the case for an LMXB. The archival NIR maps studied by \\cite{rojas11} show that no source was detected within either the \\chandra  or the ATCA radio error circles in April-October 2010 maps, down to a limit of K$_\\mathrm{s}$$\\sim$18.0. Using the relation of \\cite{prehdel95}, a visual extinction of A$_\\mathrm{V}$=39.1 magnitudes is derived from $N_\\mathrm{H}$=7$\\mathrm{\\times 10^{22}~cm^{-2}}$ (Table~\\ref{tab:efits}). Assuming a spectral type supergiant Ib (B0.5) with M$_\\mathrm{V}$ = $-$6.8 and $\\mathrm{V}$-$\\mathrm{K} = $-$0.75$ \\citep{ducati01} (i.e., an HMXB), a distance of 8\\,kpc, and using the extinction law by \\cite{cardelli89}, A$_\\mathrm{K}$/A$_\\mathrm{V}$=0.114, we obtain an estimated observed K=12.92 that would have been detected in the maps by \\cite{rojas11} (the NIR emission from an HMXB does not vary dramatically due to its bright stellar emission). Hence in the case of an HMXB hosting a supergiant, to reach the K$\\sim$18.0 limit we would need to place the source at about 83\\,kpc or more. Given our un-absorbed X-ray flux (F$_{\\mathrm{3-200\\,keV}}$=1$\\times \\mathrm{10^{-9}~erg~cm^{-2}~s^{-1}}$), this would result in an X-ray luminosity of about L$_\\mathrm{(3-200\\,keV)}$$\\sim$8$\\times$10$^\\mathrm{38}$ erg~s$^\\mathrm{-1}$ from an HMXB placed outside the Galaxy. This scenario, though not impossible, is highly unlikely since such luminosities are too high for HMXBs with the compact object accreting stellar wind.  \\\\ On the other hand, Be stars (representing the majority of companions in HMXB systems) are main sequence stars of spectral type between 09 and B2, with a spread in absolute magnitude of nearly 1 mag \\citep[see, e.g.,][]{allen84}. Be X-ray binaries, although dimmer than supergiants, are usually bright NIR sources. The Be phenomenon itself, with the presence of a decretion disk due to fast stellar rotation close to disruption, increases the luminosity of the source, at a level of about K=0.25 mag \\citep{dougherty94}. To give an example, a normal B0V star (absolute magnitude M$_\\mathrm{V}$=$-$4.1) with apparent K$\\sim$18 mag and A$_\\mathrm{V}$ = 39.1 mag would be at nearly 25\\,kpc, and the disc emission would have the effect of increasing its distance of nearly 3\\,kpc, placing this star at nearly 28\\,kpc. This is a substantial distance for instance in comparison with the faintest ($K_\\mathrm{s}$=15.75) and furthest (> 8.5\\,kpc) known candidate BeXRB in the Galaxy, described in \\cite{zurita08}. In general however, distances of HMXBs, both hosting Be and supergiant stars, are usually of the order of a few kpc, due to their lower intrinsic X-ray luminosities than LMXBs, and they are distributed mainly towards tangential directions of the Galactic arms \\citep{grimm02}.\\\\ Furthermore, direct accretion through stellar wind --- common in HMXBs --- should prevent the formation of strong and coherent jets, while the detection of radio emission from \\mysou~suggests the presence of ejection phenomena, and therefore an LMXB nature filling its Roche Lobe and forming an accretion disk.\\footnote{One could suppose the case of Cyg X$-$1, which exhibits both an accretion disk and stellar wind, but this source is the only example known in our Galaxy of BH with Roche-lobe overflow in a supergiant X-ray binary. In general, HMXBs strictly filling their Roche-Lobe are likely not observable, since the mass transfer is highly unstable and the accretion should only last for a few thousand years. Instead, there are HMXBs exhibiting Beginning Atmospheric Roche Lobe Overflow, where the massive star does not fill its Roche Lobe, but the stellar wind follows the Lagrange equipotentials, and accumulates, forming an accretion disk \\citep{bhattacharya1991}. This situation is more stable, but still rare due to the required configuration of stellar radius, orbital distance and mass ratio, since we know only 3 such systems in total, hosting neutron stars: LMC~X$-$4, Cen~X$-$3 and SMC~X$-$1.}\\\\  All these arguments suggest that the HMXB nature can be reasonably ruled out, pointing towards an LMXB nature for \\mysou. Indeed, assuming a spectral type main sequence K5 star with M$_\\mathrm{V}$=7.3 and $\\mathrm{V}$-$\\mathrm{K}$=2.66 mag \\citep{ducati01} (i.e., an LMXB) and a distance of 8\\,kpc, we would obtain K$\\sim$23 mag, compatible with the non-detection in \\cite{rojas11}, and an X-ray luminosity of about L$_\\mathrm{(3-200\\,keV)}$=7$\\times$10$^\\mathrm{36}$ erg~s$^\\mathrm{-1}$ that, though dim, is not  exceptional for the LMXB outburst luminosities.  \\subsection{A low-mass X-ray binary: geometry and compact object}\\label{sec:radio} \\begin{figure} \\includegraphics[width=0.9\\linewidth]{fig8.ps} \\caption{Radio to X-ray luminosity for a sample of black hole and neutron star LMXBs. Adapted from \\cite{coriat11}, see text for details.}\\label{fig:radiotoX} \\end{figure}  Adding the local and Galactic (ISM) absorbing column densities of Table~1, we observe that we obtain a column density up to about ten times more  than the Galactic average value expected in the source direction, $\\sim$1.3$\\times$$10^\\mathrm{22}$\\,cm$^\\mathrm{-2}$ \\citep{kalberla05}. An LMXB does not have an important wind from the companion, that is generally an old  K-M type star, but if it is seen at a high inclination (> 60$^{\\circ}$) it will appear as heavily absorbed due to the material  that gets into the line of sight (disc, blob of impact of the accreting matter onto the disc, the companion itself, etc). This is the case for the dipping, eclipsing or so-called ''coronal'' sources, i.e., sources that exhibit dips, total or partial eclipses in the soft X-ray spectra \\citep[e.g. X1822$-$371, EXO~0748$-$676, 4U~1624$-$490;][]{parmar86,white82,xiang2009,diaz06}. \\\\ In our case, it is difficult to state if we are in the ''pure-dipper'' scenario (inclination between 60 and 75 degrees) or at higher inclinations (eclipsing or coronal source).  We do not see any total eclipse in our data (zero flux) but our soft X-ray observations are either continuous 20\\,ks (5.5\\,hrs, \\chandra) or quick snapshots that add up to about 15\\,ks (\\swift) which is a small part of a typical LMXB orbital period (typically below 15--20\\,hrs). For such a dim source, we would  need to integrate many orbital periods to observe a  modulation  (dip, eclipse etc), but what is clear is that the environment we are observing is not ''clean'' (high $N_\\mathrm{H}$) and it is irregular and patchy (the three rate levels in \\chandra are not contiguous in time). \\\\ We note that the non-variability in the IBIS/ISGRI spectrum together with the absence of total eclipses (albeit the short coverage) could suggest that the Comptonizing region (responsible for the photons above 20\\,keV) would be extended and wider than the blocking material. Such ''coronal'' sources are usually known as Accretion Disc Corona sources (ADC), high inclination systems   where the accretion disc completely blocks the line of sight to the central source at all orbital phases, but that can still  be observed because the corona allows X-rays to be scattered into the line of sight. If this is the case for \\mysou, then the ''true'' X-ray flux could be higher. This is very similar to what obtained by \\cite{xiang2009} on the high inclination LMXB,  so-called ''Big Dipper'', 4U~1624$-$490, based on a \\chandra grating  observation over the $\\sim$76\\,ks binary orbit of the source. The continuum spectrum could be modeled using a single $\\Gamma$$\\sim$2 power-law partially covered by a local absorber of column density $N_\\mathrm{H}$$\\sim$8$\\mathrm{\\times 10^{22}~cm^{-2}}$ besides the Galactic one, similarly to \\mysou. Unlike \\cite{xiang2009}, our data do not allow us to constrain discrete lines in the spectrum (due to the short exposure and dim source), but we can speculate that the scenario for \\mysou~could be similar to the Big Dipper case for which \\cite{xiang2009} have shown that  X-ray variations are predominantly driven by changes in obscuration, rather than intrinsic variation of the components.\\\\   LMXBs are  known to produce radio emission, be they hosting a black hole or a neutron star. Though black hole binaries are known to be more radio-loud, there are clearly some broad similarities in their X-radio coupling, such as the association between X-ray spectral states and the presence of radio emission \\citep{fender06,corbel03,coriat11,migliari06,paizis06}. Radio emission can result either from a powerful compact jet in the hard state, or from relativistic discrete ejections at  spectral state transitions \\citep[from hard intermediate to soft-intermedite state, e.g.,][]{fender06}. While  discrete radio ejections show a radio spectrum with spectral index $\\alpha<0$,  compact jets have a flat or slightly inverted spectrum (i.e. $\\alpha\\gtrsim0$). In our case, the poorly constrained  $\\alpha$=$-0.37\\pm$0.79 does not allow us to firmly establish the nature of the radio emission, however  since \\mysou~ has not shown any important spectral transition but stayed in the hard state, it is likely that the radio emission is  associated with compact jets.\\\\ The quasi-simultaneous observation and detection of \\mysou~with \\chandra and ATCA is very important because besides leading to a constrained source position, it basically rules out the AGN and HMXB nature of the source. Indeed \\mysou~fits well in the LMXB scenario also from the radio emission point of view (regardless of the inclination). Figure~\\ref{fig:radiotoX}, adapted from \\cite{coriat11}, shows the radio luminosity against the 3--9\\,keV un-absorbed luminosity for a sample of black hole candidate LMXBs in the hard state (H~1743$-$322, GX~339$-$4 and V404~Cyg) and atoll neutron star LMXBs in the island state (Aql~X$-$1 and 4U~1728$-$34). The line is the fit to GX~339$-$4 from Corbel et al. (in preparation). \\mysou~is also shown in the plot. The assumed distance is 8\\,kpc with an un-absorbed X-ray flux of F$_\\mathrm{(3-9\\,keV)}$=7.5$\\mathrm{\\times 10^{-11}~erg~cm^{-2}~s^{-1}}$ (obtained from the least absorbed case) and a radio flux of F$_\\mathrm{(radio)}$=0.2\\,mJy.  \\mysou~seems to follow the behavior of the black hole candidate LMXB H~1743$-$322 at low fluxes, when it starts to join the standard correlation of GX~339$-$4 \\citep{corbel11,coriat11}. This could possibly suggest that \\mysou, that appears also to be too radio-bright to be a neutron star LMXB, is a black hole candidate in the hard state. We note however that we cannot firmly exclude the neutron star nature of the compact object based on the location of \\mysou~in Fig.~\\ref{fig:radiotoX} alone, since in the aforementioned ADC scenario the observed flux (hence luminosity) could be underestimated. A higher ''true'' X-ray flux would indeed move the source towards the neutron star populated area."
    },
    "1107/1107.2668_arXiv.txt": {
        "abstract": " ",
        "introduction": "} The Andromeda galaxy (M31) is an exquisite laboratory for studying stellar clusters.  Its proximity \\citep[785 kpc;][]{McConnachie05} allows for the detailed study of individual stars in clusters, while simultaneously providing a large, galaxy-wide sample of objects.  Studies of M31's stellar cluster system have been on-going since the work of \\citet{Hubble32}.  The Panchromatic Hubble Andromeda Treasury (PHAT) survey, described in Sec.~\\ref{phat}, represents a significant step forward in the study of Andromeda's cluster system, extending the sample of known clusters well into the intermediate and low mass regimes.  This dataset will inform our understanding of cluster evolution as a whole, through its wide sampling of age, mass, and galactic environment parameter space.  The survey's cluster science goals include placing constraints on cluster disruption behavior, assessing environmental dependencies of cluster formation and destruction, and measuring the star cluster initial mass function, among many others.  The first step to achieving these goals lies in the accurate identification and characterization of M31's cluster system.  We describe our current progress on this task in Sec.~\\ref{clusterwork}, and direct the reader to our forthcoming paper (Johnson et al., in prep.) for complete details. ",
        "conclusions": ""
    },
    "1107/1107.4458_arXiv.txt": {
        "abstract": "The beaming effect (aka Doppler boosting) induces a variation in the observed flux of a luminous object, following its observed radial velocity variation. We describe a photometric signal induced by the beaming effect during eclipse of  binary systems, where the stellar components are late type Sun-like stars. The shape of this signal is sensitive to the angle between the eclipsed star's spin axis and the orbital angular momentum axis, thereby allowing its measurement. We show that during eclipse there are in fact two effects, superimposed on the known eclipse light curve. One effect is produced by the rotation of the eclipsed star, and is the photometric analog of the spectroscopic Rossiter-McLaughlin effect, thereby it contains information about the sky-projected spin-orbit angle. The other effect is produced by the varying weighted difference, during eclipse, between the beaming signals of the two stars. We give approximated analytic expressions for the amplitudes of the two effects, and present a numerical simulation where we show the light curves for the two effects for various orbital orientations, for a low mass ratio stellar eclipsing binary system. We show that although the overall signal is small, it can be detected in the primary eclipse when using \\ik Long Cadence data of bright systems accumulated over the mission lifetime. ",
        "introduction": "\\label{sec:intro}  The angle between the spin axis of each star in a binary system and the orbital angular momentum axis is an important clue for studying stellar binary formation and the effect of tidal processes on  their evolution. Naively one would expect the spins of the two stars and the orbital angular momentum to be well aligned, as they all originate from the angular momentum of the same primordial molecular cloud. However, this simple view might be misleading. For example, binary stars with high obliquity could result from elongated primordial clouds \\citep{bonnell92}, or chaotic accretion during star formation \\citep{bate10}. This could also be the result of interaction with a third body through the Mazeh-Shaham effect \\citep{mazeh79}, which is essentially a combination of the so-called Kozai oscillations \\citep{kozai62} followed by tidal dissipation \\citep{fabrycky07}. A large sample of binary systems with measured obliquities is needed for studying the above processes.  To date, spin-orbit alignment has been measured, with spectroscopy, in only three eclipsing stellar binaries, two of which are aligned (V1143 Cyg, \\citealt{albrecht07}; NY Cep, \\citealt{albrecht11}), and the third misaligned (DI Her, \\citealt{albrecht09}). It is interesting to note that for transiting planetary systems the spin-orbit angle is being measured for an increasing number of systems \\cite[e.g.,][]{queloz00, winn05, triaud10, winn11}, currently including one third of the $\\sim$100 known transiting systems. Almost half were identified as not well aligned, thereby challenging the simplistic understanding of the formation and evolution of planetary systems (e.g., \\citealt{fabrycky09, triaud10, winn10}; see also \\citealt{schlaufman10}).  We show here that the beaming effect (aka Doppler boosting; e.g., \\citealt{zucker07, mazeh10, shporer10}), combined with the Rossiter-McLaughlin effect (hereafter RM; \\citealt{rossiter24, mclaughlin24}; see also \\citealt{holt1893}), induces a photometric signal during eclipse, where the shape of this signal is sensitive to the spin-orbit angle. Therefore, correctly modeling the light curve allows to measure this angle. To the best of our knowledge this approach was first used by \\cite{hills74}, for a white dwarf. More recently, this approach was noted by \\cite{loeb05} for transiting planetary systems, and by \\cite{vankerkwijk10} for eclipsing white dwarfs. This signal was theoretically modeled by \\cite{groot11} for three kinds of systems: eclipsing double white dwarfs \\citep{steinfadt10}, eclipsing O type binaries, and transiting planets orbiting an early type star. Here we present a further, fourth case: stellar binaries whose components are Sun-like stars, of spectral type F6 and later. Those stars are fundamentally different than earlier type stars and white dwarfs, as they are slow rotaters.  In fact, Sun-like stars are the primary targets of space missions conducting planetary transit surveys like CoRoT and \\ikp, where a large number of eclipsing binary systems are monitored. We discuss the possibility of applying this method to bright eclipsing binaries monitored by \\ikp, specifically the primary eclipse of low mass ratio binaries, where the low-amplitude signal can be detected by using data obtained during the lifetime of the mission.  We also show that there are in fact {\\it two} photometric signals induced by the beaming effect during eclipse, one sensitive to the star's self rotation and spin-orbit angle, while the other is not. Both effects need to be accounted for in order to correctly interpret the light curve and estimate $\\lambda$.  We describe the two effects in \\secr{effects} and give approximated analytic expressions for their amplitudes. In \\secr{sim} we present a numerical simulation of an eclipsing binary composed of an F and K type stars. We discuss the results of the numerical simulation in \\secr{dis}, including the possibility of detection with \\ik data, and possible degeneracies in the model. We summarize our work in \\secr{sum}. ",
        "conclusions": "\\label{sec:dis}  For the edge-on and aligned system presented in \\figr{F8K0}, the PRM amplitude is 55 ppm for the primary and 2.2 ppm for the secondary eclipses. These amplitudes are smaller than predicted by \\eqr{aprm} due to dilution from the eclipsing star and also because that equation was derived by assuming the eclipsing object is much smaller than the eclipsed one. When this is not the case the integration across a significant fraction of the eclipsed stars' surface acts to decrease $A_{PRM}$ by averaging down the variation in the radial component of the rotational velocity. The InOrB effect amplitude for the primary is 2.7 ppm, which is 5\\% of the PRM amplitude, not far from the 0.1 ratio predicted by \\eqr{ratio}. For the secondary eclipse the numerically measured InOrB amplitude is 6.2 ppm, almost 3 times larger than the PRM amplitude, much larger than the ratio of 0.6 predicted by \\eqr{ratio}. This is not surprising since the assumptions done when deriving \\eqr{ratio} are invalid for the secondary eclipse.  Both beaming-induced eclipse signals presented here are small, and can be detected only with ultra precise photometry. To estimate whether they can be detected by \\ik we assume an observed brightness of 12th magnitude, where according to the \\ik Mission specifications \\citep[e.g.,][]{borucki10, koch10} the photometric precision is 20 ppm over 6.5 hours (to allow detection of the 13 hours long and 84 ppm deep transit of an Earth-like planet across a Sun-like star). Assuming white noise this translates to 72 ppm per 0.5 hour. Such a precision will of course not allow a detection from a single eclipse event even for the primary eclipse. Fortunately, \\ikp's long lifetime, of 3.5 years, will allow to fold together many eclipses and decrease the scatter by $\\sqrt{N}$ where $N$ is the number of observed eclipse events. For the 10 day system adopted here it is fair to assume that 100 eclipse events will be observed, even when accounting for gaps following data download time or  spacecraft Safe Modes. Therefore, an accuracy of about 7 ppm per half an hour is predicted to be achieved over the lifetime of the mission. Such a precision should be sufficient to detect the PRM effect during primary eclipse for a system similar to the one simulated here.  Most of the \\ik targets are observed in the default Long Cadence mode. In this mode every 270 individual consecutive exposures are stacked on-board, each of 6.02 s exposure time and 0.52 s readout time. This results in an effective cycle time of 29.4 min where photons are collected during 91.4\\% of that time. The high duty cycle allows us to treat each cycle as half an hour exposure. Hence, the individual eclipse events of the system shown in \\figr{F8K0} are sampled with 12--13 exposures, which is sufficient for resolving the temporal flux variation although the variation across half an hour is averaged out. The smaller PRM signal during secondary eclipse makes it more difficult to detect, and probably not detectable with \\ikp, although the amplitude increases for stars rotating more rapidly, which is expected for younger stars due to the Skumanich spin-down law \\citep{skumanich72}.  The system we examined here has stellar components which are significantly different from each other. If we take them to be more similar, then the secondary eclipse signal becomes larger, although the primary smaller. For a system composed of two G0V stars with a self rotation velocity of 10 \\kms, our numerical calculations show that the PRM amplitude will be close to 15 ppm. For such a system the InOrB amplitude will be 11 ppm, larger than for the case of the F8V+K0V system.  The signal to noise ratio of the in-eclipse light curve is determined by the total amount of time the system spends in eclipse. Since for given stellar radii this scales with $P^{-2/3}$, eclipsing systems orbiting with a shorter orbital period will require data taken over a shorter time span in order to reach the required precision. Taking the same 12th magnitude system adopted here, only with $P_{orb}$=1.5 day gives a precision of 7 ppm per half an hour after 150 days of data, within two \\ik quarters\\footnote{\\ik quarters have a duration of about 90 days.}. However, the eclipse duration of such a short period system is only 3.07 hours. This means that while $A_{PRM}$ remains the same, each eclipse will be covered by 6--7 Long Cadence exposures, and each half of the eclipse by only 3--4 exposures. Moreover, when all Long Cadence measurements are folded together and binned the effective averaging of the flux variability goes beyond half an hour. Therefore, using Long Cadence data to detect the PRM signal of short period eclipsing binaries will be difficult, and even more so for distinguishing between the different possible spin-orbit orientations. Using \\ik Short Cadence data, with an effective integration time of 1 minute (9 readouts), could make it possible to detect the PRM effect in short period binaries. However, the Short Cadence mode is applied only to a small sample of 512 objects, about 0.3\\% of the \\ik targets, which are carefully chosen and observed in this mode usually during only one \\ik quarter.  Measuring spin-orbit alignment for longer period systems, with a period of 10 days, as simulated here, is also expected to be more useful to study binary formation, compared to short period systems. Processes resulting from tidal interaction between the two stars, including orbital circularization, synchronization between the orbital period and the stars'  rotation, and spin-orbit alignment \\citep{mazeh08}, have a time scale that is strongly dependent on the scaled semi-major axis, $a/R_p$ \\cite[e.g.,][]{zahn89}. The spin-orbit alignment specifically has a relatively short time scale, two or three orders of magnitude shorter than the circularization time scale since it involves a transfer of a smaller amount of angular momentum \\citep{mazeh08}. Therefore, it is expected that for short period binaries the spin axes of the two stars had already aligned with the orbital momentum axis, and any initial missalignment --- a relic from their formation history --- had already faded away. Since for longer period systems, with a wider orbit, this process is slower, they are better targets for studying the initial alignment and the systems' formation process. For example, the missaligned system DI Her (\\citealt{albrecht09}, see also \\citealt{mazeh08}) has an orbital period of 10.6 day.  We used the \\ik eclipsing binary catalog of \\cite{prsa11}, including 1879 systems, to estimate the number of eclipsing binary systems whose photometry will be accurate enough to allow looking for the PRM signal. We looked for systems whose eclipse duration is longer than 5 hours, and that given their orbital period and brightness (\\ik magnitude) the photometric S/N will be no less than half that of the 10 days and 12th mag system examined here. This includes 10 day systems whose brightness is down to 13.5 mag, or 12 mag systems with period up to 40 days. This cut left about 320 systems. When we restrict ourselves to low mass ratio systems, similar to the one we examined here, this number is about 50. A detection of the PRM effect and a measurement of the spin-orbit angle for the primary in even some of those systems will greatly extend the sample of eclipsing binaries for which this angle is known.  In our numerical analysis we ignored several low-amplitude effects, including ellipsoidal distortion due to tidal forces between the two stars and reflection of light, or heating of one star by the other. For a 10 day period binary both these effects are expected to have an amplitude along the orbit which is smaller than that of the orbital beaming effect \\citep{zucker07}. In addition, during eclipses the ellipsoidal and reflection effects are at extrema, meaning their induced flux variation is minimal, while the orbital beaming flux is at its phase of fastest variation. The light curves of the ellipsoidal and reflection effects during eclipse are expected to be symmetric about the mid eclipse time, and, as for the orbital beaming, they can be predicted using modeling of the out of eclipse light curve \\citep{faigler11}.  We note that we ignore here the possibility of other sources of low level noise, such as stellar noise, meaning that this measurement requires quiet stars. In case the eclipsed star shows slight activity in the form of small spots on its surface, this can lead to sinusoidal-like flux variability at the rotational period. However, for the type of system simulated here the orbital and self rotation periods are not necessarily synchronized, so the spots-induced flux variability will be averaged out once folded on the orbital period, which is determined to high accuracy by the measurement of eclipses timing.  Our method has several possible degeneracies, which if not carefully considered can lead to incorrect interpretation of the light curve, and a biased value for $\\lambda$. Significant eccentricity can produce asymmetric eclipse light curves, depending on the argument of periastron, $\\omega$ (see \\citealt{kipping08} for a description of the impact of eccentricity on planetary transit light curves). Although, eccentricity can be estimated from the orbital beaming light curve, and, the value of $e\\cos \\omega$ is accurately determined from the time difference between the primary and secondary eclipses.  A second degeneracy results from the PRM signal being mimicked by a small shift in mid eclipse time, which may give the false impression that the PRM signal does not exist in the light curve, suggesting a scenario where $i$ = 90 deg and $\\lambda$ = 90 deg, shown in \\figr{lambda} left panel. We looked into this possibility more closely by subtracting the eclipse light curve including a small shift in mid eclipse time, and without the PRM effect, from the light curve including the PRM signal and no time shift. The mid eclipse time shifts we examined were up to a few seconds, in sub second steps. We found that the peak to peak amplitudes of the difference light curves were at least 60 ppm, and the smallest difference is for time shifts close to 1 s. We expect \\ik photometry obtained during the time span discussed here to allow mid eclipse measurements of better than 1 s precision. Therefore, a solution which does not include the PRM effect  (or assumes $i$ = 90 deg and $\\lambda$ = 90 deg) and as a result gives a wrong mid eclipse time due to the above mentioned degeneracy will have a significantly lower likelihood than a solution that does include the PRM effect.  Modeling the PRM effect adds two parameters to the model, $V_{rot}\\sin I$ and $\\lambda$ of the eclipsed star. An independent estimate of $V_{rot}\\sin I$ can be achieved from a spectrum of the star, thereby constraining this parameter so the information in the PRM signal is used for fitting only a single parameter ($\\lambda$). Moreover, an independent estimate of $V_{rot}\\sin I$ is important for lifting a third degeneracy in our model, between $\\lambda$ and $V_{rot}\\sin I$, for edge-on orbits \\citep[e.g.,][]{gaudi07}.  Another possible degeneracy may result from the InOrB signal being opposite in sign to the PRM signal (see Figures~\\ref{fig:F8K0} and \\ref{fig:lambda}), giving a false impression that the PRM signal is smaller in amplitude. In general, for non-aligned and non edge-on orientations, the combined light curve shape could be different than the PRM light curve, possibly leading to a biased estimate of $\\lambda$, even though for the primary eclipse shown in Figures~\\ref{fig:F8K0} and \\ref{fig:lambda} the PRM signal is clearly dominant. However, as already noted, the InOrB light curve shape is insensitive to the spin-orbit angle or the star's rotation velocity. The only parameter to which both signals are sensitive to is the radii ratio $r$, which is accurately determined from the eclipse light curve even without accounting for the beaming-induced effects. Therefore, the InOrB light curve is determined independently of the PRM light curve, eliminating this possible degeneracy.  Lastly, we note that the oblateness of stars, due to self rotation, may also affect eclipse light curves, especially when oblateness leads to significant gravity darkening. \\cite{barnes09} showed how this will distort the shape of planetary transit light curves when the host star is a rapidly rotating early type star. However, since Sun-like stars are slow rotators this is expected to have a minor impact on the light curves (for the Sun for example the measured oblateness\\footnote{The fractional difference in the radius at the equator and the poles.} is 10$^{-5}$).  It is interesting to consider the application of the method presented here to a transiting planet orbiting a Sun-like star. However, as seen from \\eqr{aprm}, in this case $A_{PRM}$ is expected to be at the 1 ppm level, making it difficult, if at all possible to detect. The case of a transiting planet orbiting an A type star is presented by \\cite{groot11}, where the assumed rotational velocities are at the order of 100 \\kms, and the presented light curve has an amplitude of about 10 ppm."
    },
    "1107/1107.3564_arXiv.txt": {
        "abstract": "Active galactic nuclei (AGNs) with double-peaked \\OIII\\ lines are suspected to be sub-kpc or kpc-scale binary AGNs. However, pure gas kinematics can produce the same double-peaked line profile in spatially integrated spectra. Here we combine integral-field spectroscopy and high-resolution imaging of 42 double-peaked \\OIII\\ AGNs from the Sloan Digital Sky Survey to investigate the constituents of the population. We find two binary AGNs where the line-splitting is driven by the orbital motion of the merging nuclei. Such objects account for only $\\sim$2\\% of the double-peaked AGNs. Almost all ($\\sim$98\\%) of the double-peaked AGNs were selected because of gas kinematics; and half of those show spatially resolved narrow-line regions that extend 4$-$20 kpc from the nuclei. Serendipitously, we find two spectrally {\\em unresolved} binary AGNs where gas kinematics produced the double-peaked \\OIII\\ lines. The relatively frequent serendipitous discoveries indicate that only $\\sim$1\\% of binary AGNs would appear double-peaked in Sloan spectra and $2.2_{-0.8}^{+2.5}$\\% of all Sloan AGNs are binary AGNs. Therefore, the double-peaked sample does not offer much advantage over any other AGN samples in finding binary AGNs. The binary AGN fraction implies an elevated AGN duty cycle ($8_{-3}^{+8}$\\%), suggesting galaxy interactions enhance nuclear accretion. We illustrate that integral-field spectroscopy is crucial for identifying binary AGNs: several objects previously classified as ``binary AGNs\" with long-slit spectra are most likely single AGNs with extended narrow-line regions. The formation of extended narrow-line regions driven by radiation pressure is also discussed. ",
        "introduction": "\\label{sec:introduction}  The ubiquity of binary supermassive black holes (SMBHs) is expected because of two commonly accepted facts \\citep{Begelman80}: that galaxies frequently merge \\citep{Toomre72} and that SMBHs occupy the centers of the majority of massive galaxies \\citep{Kormendy95,Richstone98}. If mergers trigger or enhance nuclear accretion as predicted by simulations \\citep{Barnes96}, a significant fraction of mergers should appear as binary active galactic nuclei (AGNs).  Over the last three decades, binary AGNs have been found with separations ranging from parsec to tens of kiloparsec (kpc) scales. Thanks to the Sloan Digital Sky Survey (SDSS), we now have a large sample of binary quasars with separations of tens of kpcs; and they account for $\\sim$0.1\\% of the QSO population at $z > 1$ \\citep{Djorgovski87, Djorgovski91, Kochanek99, Hennawi06, Myers07, Myers08, Green10} and $\\sim$3.6\\% of SDSS spectroscopically selected AGNs between $0.02 < z < 0.16$ \\citep{Liu11a}. On kpc-scales, we have six convincing examples, LBQS 0103-2753 \\citep{Junkkarinen01,Shields11}, NGC~6240 \\citep{Komossa03}, 3C~294 \\citep{Stockton04}, Mrk~463 \\citep{Bianchi08}, Mrk~739 \\citep{Koss11}, and SDSS\\,J150243.1$+$111557 \\citep{Fu11c}. On parsec-scales, there is only one convincing example, the radio galaxy 0402+379, for which only the Very Long Baseline Array (VLBA) can resolve the 7-parsec-separation binary \\citep{Rodriguez06}.  Serendipitously discovered binary AGNs are of course interesting objects to study, but in order to understand the effect of galaxy interactions on nuclear activities, it is crucial to estimate the frequency of binary AGNs as a function of physical separation. Systematic searches for parsec-scale binary AGNs have focused on objects with two distinct sets of AGN emission lines, either as double-peaked broad emission lines \\citep{Gaskell83,Gaskell96,Boroson09} or as narrow lines offset from the broad lines \\citep{Komossa08,Eracleous11,Tsalmantza11}. However, double-peaked broad lines are not necessarily binary AGNs---they could be equally well explained as a Keplerian accretion disk \\citep{Chen89,Eracleous97,Eracleous03}. And the narrow-line offset systems could be gravitational wave recoiled black holes \\citep{Komossa08}. Unfortunately, it is difficult to confirm or refute these binary AGN candidates, because of the extremely high spatial resolution needed to resolve the components and/or the long time span needed to monitor the systematic velocity changes due to orbital motion.  The same technique has been used to find binary AGNs with separations larger than the scale of the broad-line regions but within $\\sim$10 kpcs. Several systematic searches have compiled AGNs with double-peaked \\OIIIexpanded\\ emission lines (dpAGNs hereafter) from the DEEP2 survey \\citep{Gerke07,Comerford09a} and the SDSS \\citep{Xu09,Wang09,Liu10a,Smith10,Liu11a}. If double-peaked \\OIII\\ emission is indicative of a binary AGN, such emission most likely represents objects with transverse separations of $\\sim$100 parsecs to $\\sim$10 kpcs \\citep[e.g.,][]{Wang09}. However, double-peaked emission lines may also arise from mechanisms other than orbital motion, such as gas kinematics in the narrow-line regions \\citep[e.g.,][]{Gelderman94,Fu09,Fischer11,Shen11}, or jet-cloud interactions \\citep[e.g.,][]{Stockton07, Rosario10}.  Unlike parsec-scale binaries, it is relatively easy to confirm or refute the kpc-scale binary AGN candidates with existing facilities. \\citet{Shen11} have recently carried out an imaging and long-slit spectroscopic study of 31 SDSS dpAGNs (mostly at declinations of $< 22^\\circ$). They identified five binary AGNs, with each resolved galaxy component associated with AGN-photoionized narrow emission lines (i.e., type-2$-$type-2 pairs). However, their study was limited by the natural seeing and the spatial coverage of the long-slit. To move forward, we decided to carry out a high-resolution imaging and integral-field spectroscopic survey of dpAGNs.  Ground-based imaging aided by adaptive optics enables efficient high-spatial-resolution imaging surveys. In \\citet{Fu11a}, we published images of 50 SDSS dpAGNs from the Keck laser guide star adaptive optics system \\citep[LGSAO;][]{Wizinowich06}. With a spatial resolution of $\\sim$0.1\\arcsec, we found $\\sim$30\\% of dpAGNs at $z < 0.6$ show discernible companions. Using the same instrument, \\citet{Rosario11} found a similar result with a smaller sample. These results showed that at least 70\\% of the dpAGNs are either single AGNs or binary type-2 AGNs with coalesced host galaxies. It is worth noting that even a merging galaxy associated with a dpAGN could be ascribed to single AGNs---the double-peaked AGN emission lines could arise from a single active component which is merging with a non-active galaxy.  We have extended our LGSAO imaging to a sample doubling the size of that surveyed by \\citet{Fu11a}. In addition, we have obtained integral-field spectroscopy of 42 dpAGNs, including most of the merging systems identified with LGSAO imaging. In this paper, we analyze the combined data set to identify kpc-scale binary AGNs and to study the statistics of dpAGN types. We describe our observations in Section~\\ref{sec:obs}. In Section~\\ref{sec:class}, we classify the AGNs based on both stellar morphology from high-resolution imaging and ionized gas morphology and kinematics from integral-field spectroscopy. In Section~\\ref{sec:comparison} we compare the classifications based on previous long-slit spectroscopy and our integral-field spectroscopy. We show that several objects classified in the literature as ``binary\" are, most likely, single AGNs with extended narrow-line regions. We conclude by discussing the binary AGN fraction, the efficiency of the double-peaked selection technique, and the formation of extended narrow-line regions (Section~\\ref{sec:discuss}).  Throughout we adopt a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_\\Lambda=0.7$ and $h \\equiv$ H$_0$/100 km~s$^{-1}$~Mpc$^{-1}$ = 0.7, consistent with the maximum likelihood estimates from WMAP \\citep{Dunkley09}. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{Statistics of Double-Peaked AGN Types} \\label{sec:summary}  The majority of dpAGNs are due either to ENLRs or to small-scale gas kinematics. ENLRs account for half (21/42) of our IFS-observed dpAGNs in both merging (13/26) and isolated sources (8/16). We thus expect ENLRs around 50\\% of dpAGNs. Unresolved gas kinematics account for 45\\% (19/42) of the IFS-observed dpAGNs, which splits into 42\\% (11/26) and 50\\% (8/16) in merging and isolated AGNs, respectively. Taking into account the bias of the IFS sample, we expect 48\\% of dpAGNs are due to unresolved nuclear gas kinematics.  We found four promising binary AGNs. Only two of these have a double-peaked line profile attributable to the relative velocity difference of the merging nuclei; for the other two the double-peaked profile is a result of gas kinematics. Binary AGNs thus account for 15\\% (4/26) of the merging dpAGNs (that we have {\\em resolved}). There are seven ambiguous cases for which we cannot rule out the binary AGN scenario, because of the limited spatial resolution of the SNIFS datacubes (Table~\\ref{tab:ifu_sample}). If we assume that all of the ambiguous cases are binaries, then we get an upper limit of 42\\% (11/26) for the binary AGN fraction in merging dpAGNs.  From our high-resolution imaging survey we found that only 29\\% (31/106) of dpAGNs are mergers; and for the purpose of this discussion, let us assume that AGNs with a single stellar nucleus cannot be kpc-scale binary AGNs. Therefore, we conclude that the binary AGN fraction in dpAGNs is 4.5$-$12\\%\\footnote{The total percentage of the above three categories exceeds 100\\% because 2$-$9 of the binary AGNs also belong to the other two categories.}. If we include the ambiguous cases, our result agrees with that of \\citet{Shen11} ($\\sim$10\\%), although we have identified a different set of binary AGNs (see \\S~\\ref{sec:comparison}). In the next subsection, we discuss the binary AGN fraction in all SDSS AGNs.  \\subsection{Binary AGN Fraction and Completeness of Double-Peaked Selection} \\label{sec:bFrac}  From a Monte Carlo simulation assuming identical intrinsic relative velocities (500 \\kms) and line widths (250 \\kms) for all systems, \\citet{Shen11} estimated that roughly 20\\% of kpc-scale binary AGNs would have been selected as a double-peaked AGN from SDSS fiber spectra. In the following, we attempt to constrain the total binary fraction of all SDSS AGNs ($f_{\\rm Bin}$) and the completeness of the double-peaked selection to kpc-scale binaries ($c_{\\rm dp}$) using the four promising binary AGNs we have identified in a sample of 26 merging dpAGNs.  Because $\\sim$1\\% of AGNs are double-peaked \\citep[$f_{\\rm dp}$;][]{Wang09,Liu10a,Smith10} and $29^{+5}_{-4}$\\% of dpAGNs are resolved mergers in our 0.1\\arcsec-resolution images, the fraction of AGNs that are binaries where orbital motion has produced the line-splitting is $0.022_{-0.008}^{+0.025}$\\% (2/26$\\times$29\\%$\\times$1\\%; errors have been propagated); and it should be the product of the total binary fraction and the double-peaked selection completeness:  \\begin{equation} f_{\\rm dpBin} = f_{\\rm Bin} c_{\\rm dp} \\end{equation}  Similarly, the fraction of serendipitous binaries among AGNs where other mechanisms produced line splitting is $0.022_{-0.008}^{+0.025}$\\% (2/24$\\times$24/26$\\times$29\\%$\\times$1\\%). Assuming that binary AGN do not tend to drive physical mechanisms (such as outflows) that can produce double-peaked emission lines, the serendipitous binary fraction should be the product of the kinematic dpAGN\\footnote{These are objects that are double-peaked because of gas kinematics instead of orbital motion.} fraction in all AGNs, the total binary fraction, and the double-peaked selection incompleteness ($1-c_{\\rm dp}$):  \\begin{equation} f_{\\rm seBin} \\simeq f_{\\rm dp} f_{\\rm Bin} (1 - c_{\\rm dp}) \\end{equation}  \\noindent For simplicity, we have used the total dpAGN fraction ($f_{\\rm dp}$) in place of non-binary dpAGN fraction, because 98\\% of the dpAGNs are due to mechanisms other than binary orbital motion (\\S~\\ref{sec:summary}).  Combining the above two equations, we obtain:  \\begin{equation} f_{\\rm Bin} = f_{\\rm dpBin} + f_{\\rm seBin}/f_{\\rm dp} \\simeq 100 f_{\\rm seBin} = 2.2_{-0.8}^{+2.5}\\% \\end{equation} \\begin{equation} c_{\\rm dp} = f_{\\rm dpBin}/f_{\\rm Bin} \\simeq f_{\\rm dpBin}/(100f_{\\rm seBin}) = 1.0_{-0.5}^{+1.6}\\% \\end{equation}  Our nominal binary AGN fraction lies close to the upper limit of that of \\citet{Shen11} (0.5$-$2.5\\%), but the double-peaked selection completeness is an order-of-magnitude lower. This implies that the characteristic relative velocity of binary AGNs is lower than the 500 \\kms\\ assumed by \\citet{Shen11}.  There is a non-negligible probability that SMBHs in a merging pair are accreting simultaneously by chance, appearing as a binary AGN. If interactions do not enhance nuclear accretion, we would expect $\\lesssim$0.3\\% of AGNs are kpc-scale binaries. This ``chance\" binary fraction is simply the product of the average AGN duty cycle \\citep[$\\lesssim$1\\% for $M_{\\rm BH} \\gtrsim 3\\times10^7$ \\msun\\ at $z = 0.5$;][]{Shankar09} and the $\\sim$29\\% merger fraction of AGNs (assuming it is the same as that of the dpAGNs). We adopted the 1\\% duty cycle because the black hole masses of the type-1 dpAGNs are all greater than $3\\times10^7$ \\msun\\ with a mean value of $3\\times10^8$ \\msun\\ \\citep{Shen11a}. The observed binary fraction ($f_{\\rm Bin}$) is already five times higher than this ``chance\" binary fraction even at the lower bound of the 1$\\sigma$ confidence interval. This result provides evidence that galaxy interactions enhance AGN activities (i.e., increasing the AGN duty cycle from $\\sim$1\\% to $8_{-3}^{+8}$\\%), although a larger sample of binary AGNs and a more sophisticated calculation incorporating AGN \\OIII\\ luminosity functions and merger rates \\citep[e.g.,][]{Yu11} are needed to draw the final conclusion.  \\subsection{The Formation of Extended Narrow-Line Regions} \\label{sec:formation}  \\begin{figure} \\plotone{f6.ps} \\caption{ Momentum injection rate vs. nuclear \\OIII\\ luminosity (2\\arcsec-diameter aperture) for the 20 SNIFS-identified ENLRs (Fig.~\\ref{fig:enlr}). The top axis shows the corresponding black hole accretion rates, calculated assuming $L_{\\rm bol} = 3500 L_{\\rm [O~III]}$ \\citep{Heckman04} and a radiative efficiency $\\epsilon = 0.1$. The dotted and dashed lines indicate the momentum injection rates from radiation pressure of the current AGN luminosity at optical depths $\\tau = 1.0$ and 0.1, respectively. \\label{fig:pdot_lo3}} \\end{figure}  Although inefficient for identifying binary AGNs, the double-peaked selection technique is efficient for identifying ENLRs around AGNs (\\S~\\ref{sec:summary}): the SNIFS observations have accumulated 20 ENLRs at $z < 0.6$. Here we attempt to constrain the formation of such extended nebulae with the current data.  The mass of a fully ionized cloud is proportional to the luminosity of hydrogen recombination lines at a given electron density ($n_e$):  \\begin{eqnarray} M_{\\rm H} & = & m_p n_e V f \\nonumber \\\\ & = & m_p L_{\\rm H\\beta} / (\\alpha_{\\rm H\\beta} n_e h\\nu) \\nonumber \\\\ & = & 6.8\\times10^7~M_{\\odot} (L_{\\rm H\\beta}/10^{40}~{\\rm erg~s^{-1}}) (1~{\\rm cm^{-3}}/n_e) \\end{eqnarray}  \\noindent where $V$ is the total volume occupied by the cloud, $f$ is the filling factor, $m_p$ is the proton mass, $\\alpha_{\\rm H\\beta}$ the case-B recombination coefficient of H$\\beta$ \\citep[$3.03\\times10^{-14}$ cm$^3$~s$^{-1}$;][]{Osterbrock06}, and $h\\nu$ the energy of an H$\\beta$ photon ($4.09\\times10^{-12}$ erg). Adopting an average \\OIIIexpanded/\\Hb\\ ratio of 6 and $n_e = 1$ cm$^{-3}$, the total masses in the ENLRs range from $8\\times10^7$ to $2\\times10^{10}$ \\msun. To avoid the ``classical\" NLR in the inner part, we integrated only the pixels at distances greater than 1\\arcsec\\ from the nuclei.  Further, we can estimate the total momentum ($p_{\\rm kin} = M_H v$) of the ionized gas with the velocity maps. Finally, the size of the ENLR (\\S~\\ref{sec:size}) divided by the maximum line-of-sight velocity gives a dynamical timescale ($t = R_{NLR}/v_{max}$) between 8 and 70 Myrs, consistent with the typical AGN lifetime. Assuming the momentum has been accumulated throughout the dynamical timescale, we can estimate the rate of momentum injection ($\\dot{p}_{\\rm kin} \\simeq p_{\\rm kin}/t$), a diagnostic that can be compared with models.  It is possible that the radiation pressure from the AGN created the ENLRs, because the gas is believed to be dusty \\citep{Dopita02, Groves04a}. Our sample is a mixture of type-1 and type-2 AGNs, so we used the nuclear \\OIIIexpanded\\ luminosity to infer the AGN bolometric luminosity \\citep[$L_{\\rm bol} = 3500 L_{\\rm [O~III]}$;][]{Zakamska03, Heckman04}. The rate of momentum injection from radiation pressure is $\\dot{p} = \\tau L_{\\rm bol}/c$, where $\\tau$ is the effective optical depth and $c$ is the speed of light. Figure~\\ref{fig:pdot_lo3} compares the measured rate of momentum injection with the expected rate of injection from radiation pressure at the current AGN luminosity. It shows that radiation pressure is sufficient to drive the outflow if the ENLR is moderately optically thick to the AGN radiation.  The uncertainties in this analysis are dominated by the lack of knowledge on the mass-weighted electron density. We have adopted a density of 1 cm$^{-3}$ because of the success of the two-phase photoionization model in fitting the emission lines of quasar EELRs \\citep{Stockton02,Fu07a}. There, we have also shown that \\OII\\ or \\SII\\ doublets overestimate mass-weighted electron densities because high-density low-excitation regions dominate \\OII\\ and \\SII\\ emission. Because the momentum is inversely proportional to the density, the data points in Fig.~\\ref{fig:pdot_lo3} easily can be offset by one dex along the vertical direction, compromising the conclusion. Detailed photoionization modeling on deeper spectra is needed to better constrain the formation of the ENLRs.  \\subsection{Size$-$Luminosity Relation of Narrow-Line Regions} \\label{sec:size}  \\begin{figure} \\plotone{f7.ps} \\caption{Size-luminosity relation of dpAGN NLRs. The data points are SNIFS-observed dpAGNs. Spatially resolved NLRs (or ENLRs) are filled circles, while unresolved NLRs are open circles with downward arrows. Binary AGNs have been excluded. The dotted, dashed, and dot-dashed lines show the size-luminosity relations from \\citet{Greene11}, \\citet{Bennert02}, and \\citet{Schmitt03}, respectively. \\label{fig:size_lo3}} \\end{figure}  AGN photoionized gas extending over kpc-scales has been detected around QSOs \\citep{Stockton87} and Seyfert galaxies \\citep{Unger87,Mulchaey96} for over three decades. Inspired by the size-luminosity relation of AGN broad-line regions \\citep[e.g.,][]{Wandel99}, several more recent studies have suggested a similar correlation for the NLRs \\citep[e.g.,][]{Bennert02,Schmitt03,Greene11}. Because of the importance of this relation to our understanding of the NLR, we re-examine this relation with the SNIFS data. Integral-field spectroscopy is well-suited for this kind of study because it combines the advantages of previous long-slit spectroscopy (better surface-brightness sensitivity) and \\hst\\ narrow-band imaging (more complete spatial coverage).  We determined the NLR radius as the longest distance from the nucleus to any lenslet with valid \\OIIIexpanded\\ detections, i.e., emission-line peak-to-noise ratios greater than 2.0. Figure~\\ref{fig:size_lo3} plots the NLR radius as a function of \\OIII\\ luminosity integrated over the entire SNIFS field of view. We measured the sizes of both ENLRs and unresolved NLRs. For the latter, the measurements are considered as upper limits.  We compare the best-fit relations from previous studies and our data. Because we reached a similar depth as the long-slit study of \\citet{Greene11}, which is an order-of-magnitude deeper than those from \\hst\\ narrow-band imaging \\citep[$> 2-12\\times10^{-15}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$;][]{Bennert02}, our measurements are more consistent with the former than the latter. Considering the upper limits from the unresolved NLRs, the dispersion of this relation is at least 0.3 dex."
    },
    "1107/1107.6031_arXiv.txt": {
        "abstract": "{The metallicity and dust-to-metals ratio of the Galaxy are fundamental parameters in understanding the interstellar medium (ISM). Currently, there is still some uncertainty surrounding these parameters. In this paper, the dust-to-metals ratio in the Galaxy is determined using the photoelectric absorption of the soft X-ray afterglows of a large sample of several hundred gamma-ray bursts (GRBs) to determine the metal column density in combination with Galactic dust maps to determine the line-of-sight dust extinction through the Galaxy in the direction of the GRB.  GRB afterglows often have large extragalactic soft X-ray absorptions and therefore the GRB sample's upper-bound will define the Galactic dust-to-metals relation.  Using a two-dimensional two-sample Kolmogorov-Smirnoff test, we determine this upper-bound and so derive the dust-to-metals ratio of the Galaxy.  We find $A_V = 2.2^{+0.3}_{-0.4}\\times10^{21}$\\,cm$^{-2}$ $N_{\\rm H}$ assuming solar, Anders \\& Grevesse (1989), metallicity.  This result is consistent with previous findings using bright X-ray sources in the Galaxy.  Using the same technique but substituting the \\ion{H}{i} maps from the Leiden-Argentine-Bonn survey for the dust maps, allows us to place a limit on the metallicity in the Galaxy.  We find a metallicity consistent with the Anders \\& Grevesse (1989) solar values often used in X-ray fitting.  Based on this and previous studies, we suggest that the metallicity of a typical ISM sightline through the Galaxy is $\\sim0.25$\\,dex higher than the current best estimate of the solar metallicity.  We further show that the dust-to-gas ratio seems to be correlated with the total gas column density, and that this may be due to the metallicity gradient observed toward the Galactic centre.  Based on the non-constant nature of the dust-to-gas ratio, we propose that the dust column density, at $N_{\\rm H} = 2.2\\times10^{21}$\\,cm$^{-2}$\\,$A_V$, represents a better proxy for the soft X-ray absorption column density than \\ion{H}{i} maps.} ",
        "introduction": "The relationship between gas, metals and dust that defines the interstellar medium (ISM), plays a central role in the properties of star-formation, and in the appearance, evolution and ultimate fate of galaxies.  However the basic quantitative relationship between gas, metals and dust is still not well-defined even for our own galaxy.  Many studies have been made over the years examining the dust-to-gas ratio in the Galaxy and more recently at cosmological distances.  Two methods dominate these analyses: a) comparing hydrogen absorption to dust extinction or reddening, using the Ly$\\alpha$ line in the UV (and sometimes H$_2$ UV lines) to get the total gas column and pairs of stars to obtain the total extinction, and b) soft X-ray photoelectric absorption that measures the total metal column density in the foreground of bright X-ray sources, and converting this to a gas column assuming a metallicity, and comparing this to a dust extinction obtained from methods like the Balmer decrement, the deviation in the optical/infrared from a blackbody spectrum or even measuring the dust column using the halo made by small angle scattering of X-rays off the dust.  The first method provides a genuine gas-to-dust ratio, in the sense that it measures the bulk of the gas directly.  However, it suffers from the deficiency that it is insensitive to ionised gas and unless the molecular lines are measured, also to H$_2$, where the fraction of hydrogen in H$_2$ may be $\\sim 0.5$ for $A_V\\gtrsim0.5$ \\citep{2009ApJS..180..125R}. Furthermore it only works along relatively low extinction lines of sight ($A_V\\lesssim2-3$) since it requires spectroscopy in the UV where the extinction is far higher and stars are often too faint to observe in Ly$\\alpha$ or H$_2$.  The X-ray absorption method works to very high extinctions ($A_V\\sim30$ or more) and measures essentially all metals whether they are ionised or even in the solid phase, providing a census of the total column density in metals.  However, it is a measurement of the metal column, not the gas column since the X-ray absorption is almost insensitive to hydrogen absorption and depends only weakly on helium. Therefore the X-ray absorption requires a metallicity conversion to move from a dust-to-metals to a dust-to-gas ratio.  It is perhaps worth noting that most studies to date in the Galaxy have either focussed on relatively nearby sets of objects or provided a small number (between 3 and about 20 for the X-ray studies) of lines of sight to locations within a few kpc of the Sun. Most studies in the Galaxy have therefore not probed the ISM of the Milky Way in either a complete or unbiased way.  These studies have consistently found a dust-to-gas ratio at a level of $N_H/A_V \\sim 2\\times10^{21}$\\,cm$^{-2}$\\,mag$^{-1}$ \\citep{1978ApJ...224..132B,1981MNRAS.196..469W,1994ApJ...427..274D,1996Ap&SS.236..285R,1973A&A....26..257R,1975ApJ...198...95G,1975ApJ...198..103R,1995A&A...293..889P,2003A&A...408..581V,2009MNRAS.400.2050G,2009ApJS..180..125R}. The statistical errors quoted for some studies have been as small as a few $10^{19}$\\,cm$^{-2}$\\,mag$^{-1}$, however the variation from study to study is closer to a few parts in $10^{20}$\\,cm$^{-2}$\\,mag$^{-1}$ This discrepancy may be related to an underestimate of the uncertainties or to an intrinsic scatter in the relation.  Gamma-ray bursts (GRBs), while found at cosmological distances, are extremely bright.  In this paper we use the large homogeneous sample of \\emph{Swift} GRB X-ray afterglows to determine upper limits to the metal column densities to several hundred lines of sight through the Galaxy.  We compare these metal column densities to all-sky hydrogen and dust surveys to obtain a new dust-to-metals ratio and metallicity value for the Galaxy.  The GRB afterglows are subject to absorption by their hosts, with a minor contribution from intervening objects, so we use a 2D 2-sided Kolmogorov-Smirnov (KS) test to overcome this limitation to define the Galactic lower envelope to their absorbing column densities.  Since it is X-ray selected, the sample is not afflicted by observational bias as UV studies are.  However, its greatest benefit over previous samples is that this sample passes lines of sight at random through the Galaxy, and passes through the entire Galaxy in every direction, providing a set of sightlines less affected by the relatively local nature of some previous studies.  In the next section I describe the sample, data reduction and analysis techniques.  In section~\\ref{sec:results}, I present the results of the study.  Section~\\ref{sec:discussion} contains an analysis of the relevance of the results and a comparison to previous efforts inside and outside our Galaxy.  In section~\\ref{sec:conclusions}, I offer my conclusions.  All errors quoted are statistical uncertainties at the 68\\% confidence level for one parameter of interest unless stated otherwise. ",
        "conclusions": "\\label{sec:conclusions}  An analysis of the dust-to-metals ratio and metallicity of sightlines through the Galaxy has been presented.  The Galactic metal column densities were determined using the lower bound of the distribution of soft X-ray absorptions of the afterglows of a large sample of GRBs detected by the \\emph{Swift} satellite.  The corresponding extinction and gas column densities were found using the dust and \\ion{H}{i} maps of \\citet{1998ApJ...500..525S} and \\citet{2005A&A...440..775K} respectively. The metal to atomic hydrogen relation is well reproduced with a metallicity $\\sim1.75$ times the solar metallicity of \\citet{2009ARA&A..47..481A}.  The best-fitting relation between metal and dust column densities is $N_{H_X}/A_V = 2.2_{-0.3}^{+0.4}\\times10^{21}$\\,cm$^{-2}$\\,mag$^{-1}$ (using AG89 abundances).  Previous observations are consistent with this result, suggesting that the metallicity for a typical ISM sightline is 0.25\\,dex higher than the current best value for the solar metallicity.  It is therefore suggested that a better reproduction of the Galactic soft X-ray absorption will be provided with a metallicity $\\sim20\\%$ \\emph{higher} than the solar metallicity of AG89.  However, it is also found that a linear representation does not reproduce the gas-to-dust relationship very well.  A gas-to-dust relationship with $N_\\ion{H}{i} = 2.00\\pm0.14\\times A_V^{0.77\\pm0.07}$ provides a much better fit to the data. It is very likely that this is predominantly a metallicity-gradient effect, and it is therefore concluded that while the gas-to-dust relation may be as given above, the best proxy for the Galactic soft X-ray absorption should be given by the dust column density with a relation of $N_{\\rm H_X}/A_V = 2.2\\times10^{21}$\\,cm$^{-2}$\\,mag$^{-1}$ or $N_{\\rm H_X}/E(B-V) = 6.8\\times10^{21}$\\,cm$^{-2}$\\,mag$^{-1}$ for an R$_V=3.1$."
    },
    "1107/1107.2424_arXiv.txt": {
        "abstract": "In previous studies, we examined turbulence-flame interactions in carbon-burning thermonuclear flames in Type Ia supernovae.  In this study, we consider turbulence-flame interactions in the trailing oxygen flames.  The two aims of the paper are to examine the response of the inductive oxygen flame to intense levels of turbulence, and to explore the possibility of transition to detonation in the oxygen flame. Scaling arguments analogous to the carbon flames are presented and then compared against three-dimensional simulations for a range of Damk\\\"ohler numbers ($\\Da_{16}$) at a fixed Karlovitz number.  The simulations suggest that turbulence does not significantly affect the oxygen flame when $\\Da_{16}<1$, and the flame burns inductively some distance behind the carbon flame.  However, for $\\Da_{16}>1$, turbulence enhances heat transfer and drives the propagation of a flame that is {\\em narrower} than the corresponding inductive flame would be. Furthermore, burning under these conditions appears to occur as part of a combined carbon-oxygen turbulent flame with complex compound structure. The simulations do not appear to support the possibility of a transition to detonation in the oxygen flame, but do not preclude it either. ",
        "introduction": "\\label{sec:intro}  A major uncertainty in the modeling of Type Ia supernovae (SN Ia) is the physical process whereby a subsonic deflagration transitions to a detonation. Such a transition seems to be required by the observations \\citep{Kozma05,Hoflich95,Mazzali07,Kasen09}, at least within the context of the popular ``single degenerate'' Chandrasekhar-mass model. Previous papers \\citep[e.g.][]{Khokhlov1997,niemeyerwoosley1997,Woosley09,Aspden10} have focused on the possibility of a transition to detonation in carbon-rich material as the deflagration enters the ``distributed burning regime'' where burning becomes slow enough for turbulence to disrupt the flame, mixing hot ash and cold fuel. Other papers \\citep[e.g.][]{Plewa04} have explored the possibility that detonation may be mechanically induced by the collision of burning waves near the surface of the white dwarf. The present situation is inconclusive.  The collisions may not be strong enough to robustly cause a detonation \\citep{Roepke07} and the amount of turbulence required for a spontaneous detonation in the distributed regime is quite large \\citep{Woosley09}.  Here we consider a third possibility - that the necessary carbon detonation actually begins as an {\\sl oxygen} detonation in a hybrid flame. This possibility has been recently considered \\citep{Woosley10} in a one-dimensional study. Carbon burning produces oxygen-rich ash that still contains a large potential reservoir of nuclear energy. The oxygen ash is produced, for a given fuel density, at a constant temperature that gradually rises as a result of oxygen burning, until, finally, a silicon-rich composition is produced. As a result of turbulence, this oxygen layer, which we shall refer to as an oxygen ``flame'', is broadened and islands of nearly isothermal conditions are produced. Here we flesh out those one-dimensional results in a series of three-dimensional simulations.  In two previous three-dimensional studies, \\citet{Aspden08a,Aspden10} (henceforth Papers I and II), turbulence-flame interactions in carbon-burning flames were examined at small and large scales, respectively.  In Paper I, it was shown that once the turbulence was sufficiently strong, the mixing of fuel and heat was driven by turbulent mixing instead of thermal diffusion. This resulted in a categorically different kind of flame, which was referred to as a distributed flame. Paper II extended these small-scale studies to (more realistic) larger length scales, where scaling relations based on the theory of \\cite{Damkohler40} were predicted to reach a limiting behavior, resulting in a so-called ``$\\lambda$-flame''.  In this paper, we focus on turbulence-flame interactions in the trailing oxygen flame, which are expected to be significantly different than in the carbon flame, due to the inductive nature of the oxygen burning.  The specific question we address is whether turbulence can lead to a greatly extended oxygen-rich region that might have properties suitable for detonation. ",
        "conclusions": "\\label{sec:conclusions}  The theoretical treatment of distributed carbon-burning thermonuclear flames from \\cite{Aspden10} has been applied to the trailing oxygen flames and compared with three-dimensional simulations over a range of Damk\\\"ohler numbers.  It was shown that for $\\Da_{16}\\ltaprx1$, turbulence does not greatly alter the flame from one in which the oxygen burns purely inductively.  Since turbulence accelerates the carbon flame however, the width of the oxygen flame is enormously broader than in the laminar case. For $\\Da_{16}\\gtaprx1$, turbulence enhances heat transfer and drives flame propagations that is {\\em narrower} than the corresponding zero turbulence inductive oxygen flame.  This is somewhat counterintuitive as turbulence typically broadens interfaces rather than sharpening them.  A consequence of burning in this limit is that the oxygen can burn faster than the inductive flame speed (but is limited by the carbon flame speed of course). Therefore, the oxygen flame does not trail behind the carbon flame (at a distance equal to the post carbon-flame velocity times the oxygen burning time scale), but burns as a compound carbon-oxygen. This suggests that a single level set is a suitable flame model for the compound flame under these conditions.  Averaging the temperature field using a cubic filter suggested that the temperature variance in at the desired conditions is too high to support the potential transition to detonation in oxygen proposed in \\cite{Woosley10}. However, this does not preclude this kind of transition under different conditions, such as higher turbulence or lower densities, and it should be borne in mind that only one such event would be required, and in the star there are many realizations."
    },
    "1107/1107.5032_arXiv.txt": {
        "abstract": "In this  paper, we compare components of the horizontal flow  below the solar surface in  and around regions consisting of rotating and non-rotating sunspots. Our  analysis suggests that there is a significant variation in both components of the horizontal flow at the beginning  of sunspot rotation as compared to  the non-rotating sunspot. The flows in surrounding areas are in most cases relatively small. However, there is a significant influence of the motion on flows in an area closest to the sunspot rotation. ",
        "introduction": "Sunspots that rotate around their umbral centers are termed as rotating sunspots. In some cases, these are identified by the rotation around another sunspot within the same active region (AR). The identification of these sunspots has been made easier by advances in the spatial and temporal resolution of recent satellite-borne telescopes (e.g. \\cite[Brown \\etal\\ 2003]{Brown03}). The origin of rotational motion is believed to be due to the shear and twist in magnetic field lines or vice-versa. It is also suggested that the magnetic twist may result from large-scale flows in the solar convection zone and the photosphere or in sub-photospheric layers. In an earlier study,  \\cite[Zhao \\& Kosovichev (2003)]{zhao03} found evidence for two opposite sub-photospheric vortical flows in the depth range of 0--12 Mm around a fast rotating sunspot in AR 9114. In this paper, we extend our earlier study of the flows beneath a rotating sunspot (\\cite[Jain \\etal\\ 2010a]{jain10a}) to the surrounding regions,  and investigate how flow fields change with depth compared to a non-rotating case. ",
        "conclusions": "We have compared the horizontal flow components in and around two sample cases of rotating and non-rotating sunspots. Our analysis suggests that both $u_x$ and $u_y$ show significant variation with depth during the course of sunspot rotation as compared to the non-rotating sunspot. The maximum change in $u_x$ is found at the beginning of sunspot rotation. The flows in surrounding regions in both cases are relatively small, however there is a significant variation in regions neighboring to the rotating sunspot. In the present study, we have considered a region that includes both rotating and non-rotating sunspots, and the quiet area around them. This is mainly due to the limitation imposed by the resolution of images on the technique.   However, the high-resolution images from Helioseismic Magnetic Imager (HMI) onboard {\\it Solar Dynamics Observatory (SDO)}  will allow us to study smaller regions within an active region that will provide a deeper insight on the dynamics of subsurface layers beneath individual sunspots, e.g. using HMI test Dopplergrams, \\cite[Jain \\etal\\ (2010b)]{jain10b} have shown that the ring-diagram technique can be reliably applied to regions as small as 5$^{\\mathrm{o}}$."
    },
    "1107/1107.2081_arXiv.txt": {
        "abstract": "We describe the successful design, implementation, and operation of a 12-GHz holography system installed on the Robert C. Byrd Green Bank Telescope (GBT)\\@.  We have used a geostationary satellite beacon to construct high-resolution holographic images of the telescope mirror surface irregularities.  These images have allowed us to infer and apply improved position offsets for the 2209 actuators which control the active surface of the primary mirror, thereby achieving a dramatic reduction in the total surface error (from 390$\\,\\mu$m to ${\\sim}240\\,\\mu$m, rms).  We have also performed manual adjustments of the corner offsets for a few panels.  The expected improvement in the radiometric aperture efficiency has been rigorously modeled and confirmed at 43~GHz and 90~GHz. The improvement in the telescope beam pattern has also been measured at 11.7~GHz with greater than 60~dB of dynamic range.  Symmetric features in the beam pattern have emerged which are consistent with a repetitive pattern in the aperture due to systematic panel distortions.  By computing average images for each tier of panels from the holography images, we confirm that the magnitude and direction of the panel distortions, in response to the combination of gravity and thermal gradients, are in general agreement with finite-element model predictions.  The holography system is now fully integrated into the GBT control system, and by enabling the telescope staff to monitor the health of the individual actuators, it continues to be an essential tool to support high-frequency observations. ",
        "introduction": "The Robert C. Byrd Green Bank Telescope (GBT) of the National Radio Astronomy Observatory (NRAO) in Green Bank, WV, is a 100-meter fully-steerable single dish radio telescope operating at frequencies from 290~MHz to 100~GHz.  The GBT uses a dual-offset Gregorian design providing a fully unblocked aperture.  With a total moving mass of 7700 metric tons \\citep{Parker05}, its design and operational status as of 2008 were summarized by \\citet{Prestage09}.  The research program to further improve the high-frequency performance of the telescope is known as the Precision Telescope Control System (PTCS)\\@.  This program has fielded several successful innovations in recent years.  In addition to an advanced pointing model including real-time thermal correction terms \\citep{Constantikes07,Constantikes04}, an optical quadrant detector is now in use by observers working at the highest frequency bands, in order to measure and correct for dynamic pointing errors due to wind-induced motion of the feedarm \\citep{Ries11}. Also, efforts are well underway to implement an advanced digital servo control system to achieve subarcsecond tracking performance.  The other major goal of the PTCS work has been to improve the surface accuracy of the GBT primary mirror.   In order to achieve and maintain an accurate paraboloidal shape, the GBT primary was built with an active surface comprising 2004 panels and 2209 actuators.  The geometry of the surface is described by \\citet{Schwab90}, and the hardware and control system is described by \\citet{Lacasse98}. The actuator home positions were set originally (in Spring 2000) via optical photogrammetry of targeted points on the surface, with the telescope parked at a single elevation.  In 2002, an elevation-dependent finite element model (FEM) was incorporated into the active surface control system in order to achieve a more uniform gain as a function of elevation.  The residual errors in the FEM correction were later measured via the out-of-focus (OOF) holography technique at 43~GHz \\citep{Nikolic07a}.  The residuals were modeled using a set of elevation-dependent Zernike polynomial expansions, and the subsequent application of the associated Zernike-based corrections resulted in an essentially flat gain curve at 43~GHz \\citep{Nikolic07b}.  At this point, the measured aperture efficiency corresponded to a surface error of about 390\\,$\\mu$m rms.  The GBT monitor and control (M\\&C) system now implements these corrections automatically at the beginning of each scan during all high-frequency observations.  Since the measured panel manufacturing error was only 75\\,$\\mu$m rms, and a model of the thermal plus gravitational error of the individual panels was computed to be $<125\\,\\mu$m rms \\citep{RSI92}, the major contributor to the total surface error was believed to be errors in the actuator zero positions derived from photogrammetry.  In late Fall 2008, we installed a Ku-band (12-GHz) holography system on the telescope in order to measure the residual errors in the actuator zero positions.  By obtaining large maps of the telescope's beam pattern, the technique of ``with-phase'' holography allows one to measure the complex electric field distribution across the aperture and subsequently use the phase distribution to construct a high-resolution image of the surface irregularities \\citep[e.g.,][]{Bennett76,Scott77}.  These images can then be used to derive corrections for adjusting the panels mechanically to reduce the alignment errors.  In this paper, we describe the instrument and present the results of the observations and the surface correction strategy that we followed during 2009--2010.  We also review the models of systematic panel-scale errors and compare them to the observed tier-averaged profiles from the holography maps. Finally, we discuss the effect of thermal environment on the ultimate surface performance of the GBT. ",
        "conclusions": "We have successfully designed and installed a 12-GHz holography system on the GBT\\@.  We have performed a campaign of high-resolution holographic imaging of the telescope aperture that has enabled us to achieve a dramatic reduction in the surface error of the primary mirror.  The expected improvements in aperture efficiency have been confirmed.  The holography images also reveal the magnitude and direction of the response of the panels to different thermal conditions. The holography system is fully integrated into the GBT control system and continues to be in regular use to monitor the health of the individual actuators, which is essential to support high-frequency observations."
    },
    "1107/1107.5668_arXiv.txt": {
        "abstract": "The first molecules detected at infrared wavelengths in the ejecta of a Type II supernova, namely SN1987A, consisted of CO and SiO. Since then, confirmation of the formation of these two species in several other supernovae a few hundred days after explosion has been obtained. However, supernova environments appear to hamper the synthesis of large, complex species due to the lack of microscopically-mixed hydrogen deep in supernova cores. Because these environments also form carbon and silicate dust, it is of importance to understand the role played by molecules in the depletion of elements and how chemical species get incorporated into dust grains. In the present paper, we review our current knowledge of the molecular component of supernova ejecta, and present new trends and results on the synthesis of molecules in these harsh, explosive events. ",
        "introduction": "Type II supernovae (SNe) are the end stage of the evolution of massive stars with initial masses on the Zero Age Main Sequence greater than 8 \\Ms. These explosive events deposit a tremendous energy to the surroundings with typical explosion energy of 10$^{51}$ ergs per event. The explosion releases a burst of neutrons and synthesised elements heavier than iron through rapid neutron capture (i.e, r-process). The explosion shock wave crosses the helium core, induces a reverse shock at the interface between the heavy elements core and the hydrogen envelope of the stellar progenitor, and reprocesses the ejecta a few hours after explosion. This reverse shock triggers the onset of Raleigh-Taylor instabilities and results in macroscopic mixing between the various core layers. The nebular phase of the SN is reached after $\\sim$ 100 days when the luminosity is sustained essentially by the radioactive decay of \\Co. Assuming \\gam~full trapping, the bolometric light curve declines according to the mean lifetime of \\Co~decay into \\Fe~($t_{56} = 111.3$ days). In view of this very harsh environment, the detection of molecules in the nebular phase a few months following the explosion of a SN came as a surprise. The locus of the first detection of molecules and dust in a SN event shortly after outburst was the Type II supernova SN1987A in the Large Magellanic Cloud. Indeed, the fundamental band $\\Delta v = 1$ of carbon monoxide, CO, at 4.65 $\\mu$m was observed between day 135 and day 260 (\\cite[Catchpole et al. 1988)]{cat88}, while the CO first overtone transition $\\Delta v = 2$ in the region of 2.3 $\\mu$m was detected at day 100 after explosion \\cite[(Danziger et al. 1988)]{danz88}. Excess emission around $\\sim$ 9 $\\mu$m was attributed to the fundamental band $\\Delta v = 1$ of silicon monoxide, SiO, and was observed as early as day 160 (\\cite[Aitken et al. 1988]{ait88}). Other molecular candidates were proposed as carriers of some unidentified bands in the IR spectrum and included CS (3.88 $\\mu$m band), H3$^+$ (3.4 $\\mu$m and 3.53 $\\mu$m bands) and CO$^+$ (2.26 $\\mu$m band). However, claims for their presence were not confirmed by theoretical models.  Since then, carbon and silicon monoxide ro-vibrational emission bands have been observed in other Type II SNe as well and appear to be the only molecules detected in SN ejecta to date. This restricted number of observed species may point to the specific environment of SN ejecta characterised by harsh conditions, the presence of helium ions, and the lack of microscopically-mixed hydrogen. This hydrogen-deprived medium fosters a 'poor' chemistry. Indeed, because of the absence of radicals like hydrides (OH, CH etc.), a fast, hot chemistry that builds up larger, more complex species (e.g., hydrocarbons, water or carbon dioxide) cannot proceed, as it is the case in other circumstellar environments [e.g., AGB stellar winds \\cite[(Cherchneff 2006)]{cher06}]. The chemical species thus formed are essentially simple diatomic molecules like CO and SiO. ",
        "conclusions": "The ejecta of SNe form two types of molecular species, those directly involved in the formation processes of dust (e.g., SiO) and those only depleting gas-phase elements (e.g., CO, O$_2$).  Members of the first category are rapidly depleted in the synthesis of dust  and are thus direct tracers of the dust formation processes. Therefore, dedicated observational surveys of these specific species at various times after outburst would be extremely useful to directly assess the efficiency with which the SN ejecta produces dust. Molecules of the second kind are of great interest as well. For the present 15 \\Ms~progenitor model with solar metallicity, they represent $\\sim$ 17\\% of the ejected mass at day 1000 and include CO, O$_2$, SiS and SO. Somewhat larger mass fractions of ejected molecules were derived for zero-metallicity SNe in the early universe (Cherchneff \\& Dwek 2009). These large molecular fractions impact the gas in two major ways. Firstly, the depletion by molecules of elements entering the processes of dust synthesis takes place and needs to be considered when assessing the amount of dust formed in SNe. This is well exemplified by the formation of gas-phase S$_2$ and SiS in zone 1A, which depletes atomic sulphur, thus impacting on the formation of FeS clusters. Secondly, molecules like CO are strong coolants through their ro-vibrational transitions and their presence will affect the thermal history of the environment where they reside. The cooling provided by CO molecules in O/C-rich clumps applied to the ejecta of SN1987A was addressed by \\cite{liu95}. They considered clumps with equal masses of atomic O and C and where 1\\% of the mass was in CO.  The cooling provided by CO could decrease the gas kinetic temperature down by a factor of three compared to CO-free gas where cooling was essentially provided by the 6300 {\\AA} and 6363 {\\AA} forbidden lines of [OI]. In the present model, the CO mass corresponds to $\\sim$ 2\\% of the O-rich core mass (i.e., zones 2 and 3) at day 400, where the mass fraction reaches $\\sim 24$\\% of the CO formation zone (i.e., zones 4A and 4B) at a similar time. The cooling provided by CO molecules must thus be strong as derived by \\cite{liu95}. Such a cooling will favour more efficient molecule and dust synthesis in the ejecta dense clumps.  Species of the second category are also crucial to the study of the SN ejecta evolution at very late time, i.e., in young SN remnants like Cas A. Molecules like CO or SiS are expected to survive in the dense, cool clumps where they formed and young remnants that have not yet been shocked by the reverse shock are expected to bear such a molecular signature. These species should thus be detectable at sub-millimetre wavelengths in young SN remnants and would not only provide clues on the local kinetics and the chemistry active in clumps or filaments but also on their chemical synthesis in the prior SN stage. \\bigskip \\acknowledgement This research is supported by the Swiss National Science Foundation grant No 20GN21-128950 through the European Science Foundation Collaborative Research Project CoDustMas."
    },
    "1107/1107.2789_arXiv.txt": {
        "abstract": "We present new, original and alternative method for searching signals coded in noisy data. The method is based on the properties of random matrix eigenvalue spectra. First, we describe general ideas and support them with results of numerical simulations for basic periodic signals immersed in artificial stochastic noise. Then, the main effort is put to examine the strength of a new method in investigation of data content taken from the real astrophysical  NAUTILUS detector, searching for the presence of gravitational waves. Our method discovers some previously unknown problems with data aggregation in this experiment. We provide also the results of new method applied to the entire respond signal from ground based detectors in future experimental activities with reduced background noise level. We indicate good performance of our method what makes it a positive predictor for further applications in many areas. ",
        "introduction": "One often encounters a problem if in the given noisy experimental data some repeatable weak signal is hidden. If such signal is really very weak, even its existence in stochastic background data is difficult to confirm, not mentioning the challenge to describe its detailed properties. This problem is relevant not only in detector data analysis in many fields of physics but it is also important in other areas of human activity where data transmission, data collection and data processing is involved -- just to mention: telecommunication, electronics, computer science, genetics, acoustics, astronomy, etc.  The search for long-memory effects in the measured  background noise is essential in proper determination of the conceivable periodic or just repeatable signal immersed in this noise. Such search is also the first step to reveal and then to separate other sources of long-memory data which, if added to the measured background noise, may finally change its nature.  In this report we introduce the novel method to examine the presence of long-memory and autocorrelations in time series. The method is based on properties of random matrices (RM)  eigenvalue spectra [1-3], however, we use RM approach in the new context.  Usually, one applies RM to multidimensional time series of data, seeking for correlations between  various $1$-dimensional subseries of these time series [4-6]. This technique analyzes eigenvalues of the correlation matrix between $1$-dimensional subseries with the corresponding eigenvalues of Wishart matrix [7,8]. In our approach, we propose  to analyze eigenvalue spectra for matrix entries built entirely from increments of $1$-dimensional time series with no correspondence to Wishart correlation matrix. The construction runs as follows.  Let $\\{x_i\\}, (i=1,...,N+1)$ is the $1$-dimensional cumulated time series with very large $N$. The time series of its increments is thus $\\{\\Delta x_i\\}$ with $\\Delta x_i = x_{i+1}-x_{i}$. We divide $\\{\\Delta x_i\\}$ into $[N/L]$ non-overlapping subseries $s_k = \\{\\Delta x_{(k-1)L+1}, ..., \\Delta x_{kL}\\}$ of length $L$ each $(k=1, ..., [N/L])$. Every subseries is then renormalized according to: \\begin{equation}% \\label{eq.1} s_k\\ \\rightarrow \\hat{s}_k = \\frac{s_k}{\\sqrt{L}\\sigma_k}, \\end{equation} where $\\sigma_k$ is the standard deviation of elements in $k$-th subseries.\\\\ We build $L\\times L$ matrices from $\\hat{s}_k$ subseries treating them as subsequent matrix rows. The first $L$ subseries build  the first matrix, the  second $L$ subseries build the second matrix, etc. This way, the $(i,j)$ entry of the $n-th$ matrix  $M^{(n)}_{ij}$ ($n=1, ..., [N/L^2]$) reads in terms of time series increments as follows: \\begin{equation}% \\label{eq.2} M^{(n)}_{ij} = \\frac{\\Delta x_{[(n-1)L+i-1]L+j}}{\\sqrt{L}\\sigma_{(n-1)L+i}} \\end{equation}  Further steps rely on examination of eigenvalue spectrum properties for the ensemble of matrices built this way and on comparison these properties with some {\\it a' priori} known spectrum patterns. Any distortion of the {\\it a' priori} known long range dependence in  the background time series data will be classified as the presence of a new signal immersed in these data. The measurement of this distortion in terms of eigenvalue spectra for matrices built according to the scheme presented above, leads not only to discovery of a new immersed signal but may give us also some information on its parameters. ",
        "conclusions": "Concluding, we presented a novel approach in detection of repeatable signals in noisy data. Our analysis was focused on the real noisy data obtained from NAUTILUS gravitational wave experiment, as an example. In order to get in future unambiguous experimental evidence of gravitational waves, one has to apply independent methods, based on various theoretical philosophies. New approach we proposed is based on RM techniques. It revealed an existing subtle difference between the simulated detector noise and truly recorded background signal. This difference was not evident so far with the use of other techniques like Fourier analysis and  has to be eliminated if very weak gravitational waves signal is going to be crosschecked by independent and diversified methods of data analysis. Moreover, we provided also the eigenvalue spectrum of the entire respond signal expected from ground based detector in future experimental activities, by diminishing its background noise level. We believe that the good performance of our method in the latter case is a good predictor of further applications in many other areas as well."
    },
    "1107/1107.4488.txt": {
        "abstract": "We  study systems of close orbiting planets evolving under the influence of tidal circularization. It is supposed that a  commensurability forms   through the action of disk induced migration and orbital circularization. After the system enters an inner cavity or the disk disperses the evolution continues under the influence of tides due to the central star which induce orbital  circularization. We derive approximate analytic models  that describe  the evolution away from a general first order resonance that results from tidal circularization  in  a two planet system and which can be shown to be a direct consequence of the conservation of energy and angular momentum. We consider the situation when the system is initially very close to resonance and  also  when the system is between resonances. We also perform numerical simulations which confirm these models and then apply  them to two and four planet systems chosen to have  parameters related to  the GJ581 and HD10180 systems. We also estimate the tidal dissipation rates through effective quality factors that could result in   evolution to observed period ratios within the lifetimes of the systems. Thus the survival of, or degree of departure from,  close  commensurabilities  in   observed systems may be indicative of the effectiveness of tidal disipation,  a feature which in turn may be related to the internal structure of the planets involved. ",
        "introduction": "\\label{intro} Planetary systems  containing  hot Neptunes  and  hot super--Earths  have been observed recently. A  system   of this kind consists of the  four planets around the M--dwarf GJ~581 (Bonfils et al. 2005, Udry et al. 2007, Mayor et al. 2009a).  The projected masses of the planets are 1.9, 15.6, 5.4 and 7.1~M$_{\\oplus}$ and the periods are  3.15, 5.37, 12.93 and 66.8 days, respectively. % These corresponds to semi--major axes of 0.041, 0.07, 0.22 and 0.03~au, respectively. Other such multiple systems   are  that around  HD~40307 (Mayor et al. 2009b) which consists of  three  planets with projected masses of 4.2, 6.9 and 9.1~M$_{\\oplus}$ and periods of  4.31, 9.62 and 20.46 days, respectively % a=0.047 , 0.081 , 0.134 and that around HD 10180 (Lovis et al. 2010) which consists of seven planets with projected masses 1.35, 13.10,  11.75,  25.1,  23.9,  21.4  and    64.4~M$_{\\oplus}$ and periods of  1.18,   5.76, 16.36, 49.74, 122.76,  601.2, and   2222 days respectively. The innermost member of the latter system has yet to be confirmed.   Migration due to tidal interaction with the disk is a  possible  mechanism through  which planets end up on short  period orbits, as {\\em in situ} formation implies  very massive discs (e.g., Raymond et al. 2008). Terquem \\& Papaloizou 2007 (see also  Brunini \\& Cionco 2005)  proposed a scenario for forming hot super--Earths in which a population of cores that formed  at some distance from the central star migrated inwards due to interaction with the disk.  These collided and  merged as they  went~. This process  could produce systems of planets with masses in the earth mass range,  located inside an assumed  disk inner edge, on short period orbits with mean motions  of neighbouring planets that frequently exhibited near commensurabilities. However,  tidal circularization of the orbits induced by tidal interaction with the central star, together with later close scatterings and  mergers tended to cause the system to move away from earlier established commensurabilities  to an extent determined by the effectiveness of these processes.  Papaloizou \\&  Terquem (2010) considered  the system around HD~40307 for which the pairs consisting of the innermost and middle  planets and the middle and outermost planets are near but not very close to a  pair of 2:1 resonances. In spite of this it   was found that   secular effects produced by the action of the resonant angles coupled with the action of tides from the central star could cause the system to increasingly separate from commensurability. Resonant effects can arise even when departures from strict commensurability are apparently large because  tidal circularization produces small eccentricities which, for first order resonances,  can be consistent with resonant angle libration  (see Murray \\& Dermott 1999).  In this paper we undertake a  further study of systems of close orbiting planets evolving under the influence of tidal circularization. We present simple analytic models  describing the evolution away from a general first order resonance for a two planet system under the influence of tidal circularization, describing the situation both when the commensurabilty is very close and  also  when the system is between resonances. We also perform numerical simulations of two and four planet systems chosen to have  parameters related to  the GJ581 and HD10180 systems~. We consider the situation when various commensurabilities result through the action of assumed disk induced migration and orbital circularization rates, estimating  the magnitudes of tidal quality factors that could produce  evolution to observed period ratios within the lifetimes of the systems.   The plan of this paper is as follows. In  section\\ref{simple1} we consider  a coplanar system of  planets in near circular orbits in which orbital  energy is dissipated while its total angular momentum is conserved. The  system is expected to spread  in a  similar manner  to  a viscous accretion disk (see Lynden-Bell \\& Pringle 1974).  For a two planet system, this radial  spreading will always  lead to evolution away from  an  initially  close  commensurability.  In sections \\ref{coords}, \\ref{disk tides}  and  \\ref{p+1p}, we carry out an anlaytic study of two planets near to a first order commensurability  under the influence of tidal circularization. For planets in the mass range we consider, it is readily estimated that tides raised on the planet are very much more important than tides raised on the star (eg. Goldreich \\& Soter 1966, Barnes et al.  2009). In addition  the orbital decay timescale due to tides raised on  the star may be   estimated to be  much longer than any timescale of interest (eg. Barnes et al. 2009). Thus  tides raised on the star have  been  neglected.  We consider the initial evolution away from a  close first order commensurability in section  \\ref{tight} and go on to consider the case of evolution when the commensurability is not so close, or the system is between commensurabilities in section \\ref{loose}~. As expected from the simple arguments  given in section \\ref{simple1}, the system departs from an initially close commensurability  moving to a neighbouring one.  We go on to perform   numerical simulations of multiplanet systems  in section \\ref{Numsim}. Various commensurabilities between pairs of planets are set up by applying dissipative forces assumed to arise from a disk, that lead to orbital migration and circularization. These forces were then removed corresponding to assumptions of either entry into an inner cavity or removal of the disk. The evolution of the system under tidal circularization caused by interaction with the central star was then followed. For illustrative purposes we consider  two planet systems  with  parameters corresponding to the two innermost planets in the  GJ581 system. In section \\ref{3:2} we consider a system  that formed  a 3:2 commensurability which then evolved under orbital circularization indicating that the model system  could attain the period ratio appropriate to  the actual system  if the tidal parameter $Q'$  introduced by Goldreich \\& Soter (1966)   $ \\sim 100.$ The situation when the system began with disk parameters that led to a 5:3 commensurability is then  similarly  studied in section \\ref{5:3}. We go on to consider the effect of  adding the additional planets in the GJ581 system  in section \\ref{extraplanets}.  As there are examples of low mass planetary systems such as  HD 10180 which have separations of pairs of planets, the third and fourth innermost in that case, that indicate there may have been  a past proximity to a 3:1 commensurability, we consider  an  exploratory simulation  of a system for which the initial disk evolution  sets up  a 3:1 commensurability  in section \\ref{3:1}. Finally in section \\ref{Discuss} we summarize and discuss our results.   % % 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}   %\\begin{acknowledgements} %If you'd like to thank anyone, place your comments here %and remove the percent signs. %\\end{acknowledgements}  % BibTeX users please use one of %\\bibliographystyle{spbasic}      % basic style, author-year citations %\\bibliographystyle{spmpsci}      % mathematics and physical sciences %\\bibliographystyle{spphys}       % APS-like style for physics %\\bibliography{}   % name your BibTeX data base  % Non-BibTeX users please use %\\begin{thebibliography}{} % % and use \\bibitem to create references. Consult the Instructions % for authors for reference list style. % ",
        "conclusions": "\\label{Discuss} In this paper we have studied systems of close orbiting planets evolving under the influence of tidal circularization. We considered  the situation where  the system evolved under the influence of disk tides to form a commensurability. After the disk tides ceased to operate, either because of entry into an inner cavity, or because of  loss of the disk,  the operation of  tidal circularization caused increasing  departure from any close  commensurability as time progressed.  In  section\\ref{simple1} we pointed out that  a system of planets in near circular orbits is expected to separate on average as energy is dissipated while angular momentum is conserved. This is also expected in the very similar situation of an accretion disk evolving under a viscosity (see Lynden-Bell \\& Pringle 1974). In the simplest case of two planets, this inevitable increasing physical separation has to lead to the  increasing departure from any  initial commensurability.  In sections \\ref{coords}, \\ref{disk tides}  and  \\ref{p+1p}, we developed a formalism that could be adapted study the evolution of two planets near to a first order commensurability  under the influence of tidal circularization. This was then applied to a system with a  tight commensurability in section  \\ref{tight}. An expression for the  departure  from commensurability,  indicating this to be $\\propto t^{1/3},$ was given  (see equation (\\ref{sepeqt})). The discussion was then extended to the situation when the two planets were  not necessarily in  a close commensurability  in section \\ref{loose}. The orbital evolution of the planet in that case, leading to a neighbouring  commensurability, was then considered in section \\ref{orbevol}.  In order to confirm the analytic modeling,  numerical simulations were under taken in section \\ref{Numsim}. We were able to set up systems of low mass planets in varying commensurabilities, depending on the strengths of the disk tides leading to orbital migration and circularization, with weaker tides in general leading to more widely separated commensurabilities.  We focused on a two planet system which had the same parameters as the innermost two planets as the GJ581 system in section \\ref{3:2}.  This formed a 3:2 commensurability which then evolved under orbital circularization. This model system attained the period ratio of the actual system after $\\sim 10^8 yr.$ when $Q'\\sim 1.$ Simple extrapolation thus indicates that tidal evolution could have moved the system to the present  period ratio of $1.7$  from a 3:2 commensurability if $Q' \\sim 100.$ Similarly  the situation  when the system initially attained a 5:3 commensurability was studied in section \\ref{5:3}. In this case the evolution quickly adapted to evolve as for the case with the initial 3:2 commensurability when that had reached the same period ratio.  However, a larger value $Q'$ would suffice to  cause  the period ratio to move from  5:3  to the observed one within a given life time. The effect of  adding  the additional planets in the GJ581 system was considered in section \\ref{extraplanets}. In that case over the  long term  the extra planets did not greatly affect the evolution. However,   for a brief period the third planet moved outwards taking up the angular momentum of the innermost planet rather than the second, with the consequence that the period separation rate for the innermost pair was slowed. Thus a pair of planets may not always  evolve independently of others in the system, a feature  that requires further study.  Finally the evolution of a  system that formed  a 3:1 commensurability was considered in section \\ref{3:1}. The model system adopted had similar parameters  to the third and fourth innermost planets in the  HD 10180 system. This case required slow disk migration and weak circularization with the result that the resonance involved high eccentricities.  Because of these the commensurability could be maintained for a while under orbital circularization. However,  eventually the system increasingly   departed from it as in the other cases. Finally all of our results indicate that  if $Q'   <\\sim 100,$  commensurabilities would have been significantly affected by tidal effects related to orbital circularization. Thus the survival of close  commensurabilities in   observed systems may be indicative of the presence of large $Q'$ values, a feature which in turn may be related to the internal structure of the planets involved.   %\\end{appendix}"
    },
    "1107/1107.5817_arXiv.txt": {
        "abstract": "Cosmic ray acceleration through first-order Fermi acceleration in a collisionless plasma relies on efficient scattering off magnetic field fluctuations. Scattering is most efficient for magnetic field fluctuations with wavelengths on the order of the gyroradius of the particles. In order to determine the highest energy to which cosmic rays can be accelerated, it is important to understand the growth of the magnetic field on these large scales. We derive the growth rate of the long-wavelength fluctuations in the linear regime, using the kinetic equation coupled to Maxwell's equations for the background plasma. The instability, driven by the cosmic ray current, acts on large scales due to the stress tensor and efficient scattering on small scales, and operates for both left- and right circular polarisations. This long-wavelength instability is potentially important in determining the acceleration efficiency and maximum energy of cosmic rays around shock waves such as in supernova remnants. ",
        "introduction": "\\label{sec:intro}  The acceleration of Galactic cosmic rays up to energies of $10^{15}$~eV is widely believed to result from first order Fermi acceleration \\citep{1977Axfordetal, 1977Krymskii, 1978Bella, 1978Bellb, 1978BlandfordOstriker} in supernova remnant blast waves. This mechanism relies on a converging flow, where cosmic rays gain energy every time they cross the discontinuity, and is also referred to as diffusive shock acceleration (DSA). On either side of the shock the cosmic rays isotropise by scattering off magnetic fluctuations, while the plasma is highly collisionless. The length scale of the magnetic fluctuations and the amplitude determine the scattering efficiency. The more quickly the particles isotropise, the quicker they may diffuse back accross the shock and the shorter the acceleration time.  There exists a variety of cosmic ray streaming instabilities that amplify the local magnetic field and increase the scattering efficiency of the plasma. For a recent review, see e.g.~\\citet{2011Bykovetalreview}. On small scales, i.e.~smaller than the gyroradius, the streaming cosmic rays induce a return current that creates a force perpendicular to the background magnetic field. This results in a non-resonant instability that can amplify the magnetic field to many times its initial level in the nonlinear regime \\citep{2004Bell}. Since scattering is most efficient when the larmor radius is of the order of the wavelength of the magnetic fluctuations, it is therefore expected that this instability is effective in confining the low-energy cosmic rays.  Resonant instabilities arise in the regime where the parallel component of the wavelength corresponds to the gyroradius of the streaming cosmic rays \\citep{1967Lerche,1969KulsrudPearce,1974Wentzel,1975Skillingc}. When the perturbed field becomes of the same order as the background magnetic field ($\\delta B/B > 1$) the resonance is lost. The saturation level of the waves therefore is rather low, and this instability is not effective in amplifying the magnetic field to many times its zeroth order level.  Whether the highest-energy cosmic rays are efficiently confined in the shock region depends on whether instabilities develop on scales larger than the gyroradius of the generating cosmic rays. In this paper we show that an instability acts in this regime that arises from the cosmic ray current and is mediated by the stress tensor. The coupling of the cosmic rays to the plasma on the small scales, due to effective scattering as a result of the non-resonant instability \\citep{2004Bell}, further aids in making both left and right-handed polarisations unstable.  We use the Boltzmann transport equation together with Maxwell's equations for the background plasma, to derive the growth rate of magnetic fluctuations on the large scale. Rather than limiting ourselves to the diffusion approximation, we include higher order anisotropy and find that the stress tensor takes an active part in mediating the instability. The small-scale scattering couples the two polarised modes and integrates the short-scale with the large-scale instability. The growth rate is found to steeply depend on the wavenumber, resulting in an rapid decline with wavelength.   We limit ourselves to a mono-energetic beam of particles. This will allow us to get a cleaner picture of how exactly the instability works. The distribution function generally is expected to follow a power law beyond the thermal peak, and has a cut-off at the high-energy end. If diffusion is a function of momentum, which is generally assumed, the high-energy cosmic rays diffuse farther upstream of the shock than the lower energy ones. The low-energy cut-off therefore increases with distance upstream of the shock. Since the number of particles steeply declines with energy, it can be argued that the instability is dominated by the cosmic rays with the lowest energy in the local plasma, being a function of distance ahead of the shock.  In Section~\\ref{sec:method} we show how we derive the dispersion relation for the instability that results from including the stress tensor and scattering on small scales. In Section~\\ref{sec:k0limit} we evaluate the result and look at the growth rate of the instability for a range of collision efficiencies and wave numbers. We show the differences between inclusion or neglect of both the stress tensor and the small-scale diffusivity in Section~\\ref{sec:diff_f1f2}. We discuss consequences for confinement of cosmic rays in Section~\\ref{sec:resonance}. In Section~\\ref{sec:f3} we show that in most regimes it suffices to limit ourselves to the second order anisotropy (i.e.~the stress tensor) and illustrate the difference when higher order anisotropy of the distribution function is taken into account. Finally, we discuss its implications and conclude in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} We have derived the dispersion relation for an instability that is driven by the zeroth order cosmic ray current. The long-wavelength component is mediated by the stress tensor, and coupling between the left- and right-hand modes occurs when scattering of cosmic rays off small scale turbulence is taken into account.  The streaming of cosmic rays ahead of the shock induces a return current in the plasma that triggers instabilities from large scales down to the Alfv\\'enic regime on very small scales at which the magnetic tension becomes efficient in damping the oscillations. We did not include the Alfv\\'en term (magnetic tension), which only becomes of importance beyond $k r_g > 10^3$ for our parameters.  On small scales our analysis reproduces the hydrodynamical non-resonant instability as described in \\citet{2004Bell}. The growth rate at small scales is very high and the amplified magnetic field can reach many times its zeroth order strength. In an environment such as that upstream of a SNR blast wave, the efficient generation of magnetic field amplification on small scales is expected to increase the scattering efficiency of the lower energy cosmic rays. We find that this in turn seems to couple the left- and right-hand modes, resulting in a small left-handedly polarised contribution to the short-scale instability, in addition to the dominant right-handed mode.  On long scales, predominantly the left-hand mode is unstable. This effect only arises when anisotropy of order $\\fv_2$ (and higher) is taken into account. The long-wavelength mode behaviour already converges when orders of $\\fv_2$ are included, as we find that higher-order anisotropy terms do not alter the instability significantly. Inclusion of the collision term results in coupling between the two polarisations such that both modes are also unstable in the long-wavelength limit. The growth rate of the long-wavelength part of the instability is a function of $k^{3/2}$, which results in a rapid decline of the growth rate with wavelength. In a recent paper, \\citet{2011Bykovetal} found a larger growth rate in this regime using a different method that includes the ponderomotive force resulting from the instability on short scales.  Since we are assuming a mono-energetic beam of particles not all the effects around the resonance points are captured in this approach. However, since the number of cosmic rays steeply declines with energy, it can be expected that the instability is dominated by the cosmic rays with the lowest energy in the local plasma, being a function of distance ahead of the shock. Including higher orders of $\\fv$ shows converging behaviour in the resonance region, and the peaks that arise as a result of just including up to $\\fv_2$ smooth out.  Alternatively, the magnetic fluctuations in the resonance regime could be dominated by filaments resulting from the short-scale nonresonant instability that have grown in structure to length scales that are in resonance or beyond. It will be important to evaluate the non-linear regime and saturation levels to determine the actual efficiency of this instability in confining high energy cosmic rays.  What happens farther upstream, in the region where only the high-energy cosmic rays diffuse out to, while the lower-energy cosmic rays are confined closer to the shock region, is uncertain. If effective diffusivity only operates when a significant fraction of the low-energy cosmic rays is around, the two regimes may not couple, causing the long-wavelength instability to occur exclusively in the left-handed variety and vice versa for the short-scale instability.  An interesting result is that the dominant mode on scales larger than the gyroradius is the left-hand polarised one. This polarisation is more efficient in scattering cosmic rays since the perturbed $j \\times B$ force continues to act in the same direction over longer scales. The deflection therefore is larger, and scattering more efficient. Therefore, even though the growth rate of this mode is smaller, the long-wavelength instability may still be very effective in confining cosmic rays that have gyroradii larger than the driving cosmic rays."
    },
    "1107/1107.3436_arXiv.txt": {
        "abstract": "{ We consider the possibility to observationally differentiate the Standard Model (SM) Higgs driven inflation with non-minimal coupling to gravity from other variants of SM Higgs inflation based on the scalar field theories with non-canonical kinetic term such as Galileon-like kinetic term and kinetic term with non-minimal derivative coupling to the Einstein tensor.\\\\ In order to ensure consistent results, we study the SM Higgs inflation variants by using the same method, computing the full dynamics of the background and perturbations of the Higgs field during inflation at quantum level.\\\\ Assuming that all the SM Higgs inflation variants are consistent theories, we use the MCMC technique to derive constraints on the inflationnoary parameters and the Higgs boson mass from their fit to WMAP7+SN+BAO data set. We conclude that a combination of a Higgs mass measurement by the LHC and accurate determination by the PLANCK satellite of the spectral index of curvature perturbations and tensor-to-scalar ratio will enable to distinguish among these models. We also show that the consistency relations of the SM Higgs inflation variants are distinct enough to differentiate the models.} ",
        "introduction": "There has been much interest recently in models of inflation in which the Standard Model  (SM) Higgs boson non-minimally coupled to the  Ricci scalar can give rise to inflation without the need for additional degrees of freedom to the SM \\cite{BezSha08,BarSta08,BerSha09,deSimone09,BGS09}. This scenario is based on the observation that the problem of the very small value of the Higgs quadratic coupling required  by the Cosmic Microwave Background (CMB) anisotropy data can be solved if the Higgs inflaton has a large non-minimally coupling with gravity \\cite{BarKa94}. The resultant Higgs inflaton  effective  potential in the inflationary domain is effectively flat and can result in successful inflation for values of non-minimally coupling constant  $\\zeta \\sim 10^4$  ($\\zeta$-inflation), allowing for cosmological values for Higgs boson mass in a window in which the electroweak vacuum is stable and therefore sensitive to  the field fluctuations during the early stages of the universe \\cite{Espinosa08}. \\\\ Limits of validity of $\\zeta$-inflation have been recently debated by several authors. Specifically, Barbon \\& Espinosa (2009) argued that the large coupling of Higgs inflaton to the Ricci scalar makes this model invalid beyond the ultraviolet cutoff scale $\\Lambda_{\\zeta} \\simeq M_{P}/\\zeta$ (hereafter $M_P=2.4 \\times 10^{18}$ GeV is the reduced Planck mass) which is below the Higgs field expectation value  at $\\cal N$ {\\it e}-foldings during inflation, $\\phi \\simeq \\sqrt{N}M_P/\\sqrt{\\zeta}$. As consequence, at the ultraviolet cutoff scale  $\\Lambda_{\\zeta}$ at least one of the cross-sections of different scattering processes hits the unitarity bound \\cite{Burgess09}. The fact that the quantum corrections due to the strong coupling to gravity makes the perturbative analysis to break down at energy scales above $\\Lambda_{\\zeta}$ was interpreted as a signature of a new physics, implying higher dimensional operators at energies above $\\Lambda_{\\zeta}$. Few authors addressed the issue of  $\\zeta$-inflation naturalness with respect to perturbativity and unitarity violation in Jordan and Einstein frames \\cite{Lerner10a,Lerner10b,Burgess10,Hertzberg10}. They show that the apparent breakdown of the theory in the Jordan frame does not imply new physics, but a failure of the perturbation theory in the Jordan frame as a calculation method. These works demonstrate that, for inflation based on a single scalar field with large non-minimal coupling, the quantum corrections at high energy scales are small, making the perturbative analysis valid and  consequently there is no a breakdown of unitarity at the energy scale $\\Lambda_{\\zeta}$. \\\\ Recent works  \\cite{Bez11,Ferrara11,Atkins_a} revisit the  issue  of self-consistency of  $\\zeta$-inflation model emphasizing that the scale of unitarity violation depends on the size of the Higgs background field. They found the cutoff scale $\\Lambda_{\\zeta}$  higher than the relevant dynamical scales throughout the whole history of the universe, including the in\ufb02ationary epoch and reheating, indicating that the theory does not violate unitarity at any time.\\\\ Although the above conclusion theoretically may allow to work within the regime of validity of the effective theory for inflationary calculations, other variants of Higgs inflation within the SM have been proposed. These models are based on scalar field theories with non-canonical kinetic term \\cite{Picon} and are motivated by some particle physics models of inflation such us the Dirac-Born-Infed inflation \\cite{Ali,Easson}.  A model with non-minimal derivative couplings was proposed in \\cite{Amendola,Capo_a,Capo_b} in context of in\ufb02ationary cosmology, and recently, the inflation driven by a scalar field with a Higgs potential and non-minimal kinetic couplings to itself and to the curvature was proposed in \\cite{Granda_a,Granda_b}. In the new model of Higgs inflation (E-inflation) an enhanced kinetic term in the Lagrangian propagating no more degrees of freedom than the minimally coupled scalar field in General Relativity (GR) is obtained by considering non-minimal derivative coupling of the SM Higgs boson to the Einstein tensor \\cite{Germani10a}. The power counting analysis indicates that the scale at which the unitarity bound is violated in this model is $\\Lambda (H) \\simeq (2H^2 \\kappa)^{1/3}$, where $H$ is the Hubble expansion during inflation. Requiring the unitarity constraint $H \\ll \\Lambda$ to be satisfied, it is shown that the non-minimal derivative coupling allows a Higgs boson self-coupling value $\\lambda$ within the limits expected from the collider experiments while the cosmological perturbations in this model are  consistent with present cosmological data \\cite{Germani10b,Atkins_b}.  The Galileon models of inflation (G-inflation) are constructed by introducing a scalar field with a self-interaction whose Lagrangian is invariant under Galileon symmetry  $\\partial_{\\mu}\\phi \\rightarrow \\partial_{\\mu}\\phi +b_{\\mu}$, which maintains the equation of motion of the scalar field as a second-order differential equation, preventing the theory from exhibiting new degrees of freedom and avoiding ghost or instability problems  \\cite{Nicolis09,Deffayet09a,Deffayet09b}. Recently, an inflation model in which inflation is driven by the SM Higgs scalar field with a Galileon-like kinetic term  has been proposed \\cite{Koba10}. The dynamics of background and perturbations of G-inflaton show new features brought by the modified kinetic term in the Lagrangian, when compared with the standard slow-roll inflation, such as: violation of the standard consistency relation \\cite{Koba10,Kama10}, violation of the Null Energy Condition \\cite{Creminelli10}, generation of isocurvature perturbations and large primordial non-Gaussianity \\cite{Mizuno10,Burrage11,Felice11a,Koba11}.  In this paper we analyze to possibility to distinguish among the SM Higgs inflation variants by using the existing cosmological and astrophysical measurements. We will adopt a similar philosophy as in  \\cite{Lerner11}, considering that the SM Higgs inflation variants are consistent theories  and  therefore should be subject to rigorous tests against experimental data.\\\\ Cosmological constraints on the SM Higgs driven  $\\zeta$-inflation have been discussed by a number of authors \\cite{BarSta08,BerSha09,deSimone09}. These constraints are based on mapping between the Renormalization Group (RG) flow equations and the spectral index  of the curvature perturbations obtained by the WMAP team, parametrized in terms of the number of ${\\it e}$-foldings till the end of inflation. In a recent paper \\cite{Popa10} we studied the implications  of the full dynamics of background and perturbations of the Higgs field during $\\zeta$-inflation at quantum level and  derive constaints on the inflationary parameters and the Higgs boson mass from the analysis  of WMAP 7-year CMB measurements complemented  with astrophysical distance measurements \\cite{Komatsu10,Larson}. \\\\ Our goal in this paper is to refine the computation for $\\zeta$-inflation by considering in the analysis the improved RG flow equations, accounting at the same time for the possible uncertainties in the theoretical determination of the reheating temperature. We also make a similar analysis for  G-inflation and E-inflation models and derive constraints on the inflationary parameters and Higgs boson mass from their fit to the same data set, ensuring in this way that the differences in the predictions of the observables are due to the differences in the underlying theories of the SM Higgs inflation and therefore can be used to distinguish among them.  The paper is organized as follows. In Section~2 we introduce the there models of the SM Higgs inflation that we consider: $\\zeta$-inflation \\cite{BezSha08}, G-inflation \\cite{Koba10} and E-inflation \\cite{Germani10a,Germani10b}. In Section~3 we compute the  RG improved  Higgs  field potential and in Section~4 we present our main results. In Section~5 we draw our conclusions. \\\\ Throughout the paper we consider an homogeneous and isotropic flat background described by the Friedmann-Robertson-Walker (FRW) metric: \\begin{equation} \\label{FRW} {\\rm d}s^2=g_{\\mu,\\nu}\\,{\\rm d}x_{\\mu}\\,{\\rm d}x^{\\nu}  =  -{\\rm d}t^2+a^2(t){\\rm d}x^2 \\,, \\end{equation} where $a$ is the cosmological scale factor, $\\kappa^2 \\equiv 8 \\pi M_{pl}^{-2}$ (where $M_{pl}\\simeq 1.22 \\times 10^{19}$ GeV is the Planck mass), overdots denotes the time derivatives and $_{,X} \\equiv \\partial/ \\partial {X} $. \\\\ We also remind that the canonical tree-level action describing the SM Higgs inflation is given by: \\begin{equation} \\label{action0} S_0=\\int{{\\rm d}^4\\,x}\\sqrt{-g}\\, \\left[\\frac{1}{2\\kappa^2}{\\bf R}+ {\\cal L}_{SM}\\right]\\,, \\end{equation} where ${\\bf R}$ is the Ricci scalar and ${\\cal L}_{SM}$ is the tree-level SM Higgs Lagrangian: \\begin{equation} \\label{L_SMG} {\\cal L}_{SM}= g^{\\mu \\nu}{\\cal D}_{\\mu}{\\cal H}^{\\dagger}{\\cal D}_{\\nu}{\\cal H} -\\lambda \\left({\\cal H}^{\\dagger}{\\cal H} -v^2\\right)^2 \\,. \\end{equation} In the above expression ${\\cal H}$ is the Higgs boson doublet, ${\\cal D}_{\\mu}$ is the covariant derivative with respect to $SU(2) \\times SU(1)$ and $v$ is the vacuum expectation value (vev) of the Higgs broken phase of the SM.  Rotating the Higgs doublet such that ${\\cal H}^T=(1/\\sqrt{2})(0,v+\\phi)$ and assuming that $\\phi$ is much greather than vev during inflation we have: \\begin{eqnarray} \\label{L_SM} {\\cal L}_{SM}=P(\\phi,X)=X-V(\\phi)\\,, \\hspace{1cm} \\label{pot} X=-\\frac{1}{2}g^{\\mu\\nu}\\partial_{\\mu}\\partial_{\\nu}\\phi \\,, \\hspace{1cm} V(\\phi)=\\frac{\\lambda}{4}\\phi^4\\,, \\end{eqnarray} where $\\lambda$ is the tree-level Higgs self coupling and $V(\\phi)$ is the corresponding Higgs potential. ",
        "conclusions": "Although the SM Higgs inflaton non-minimally coupled to gravity can account for inflation ($\\zeta$-inflation), other variants of  Higgs inflation within the SM based on scalar field theories with non-canonical kinetic term such as Galileon-like kinetic term (G-inflation) and kinetic term with non-minimal derivative coupling to the Einstein tensor (E-inflation) have been proposed. In this paper we have analyzed the possibility to distinguish $\\zeta$-inflation, G-inflation and E-inflation models through their predictions for  $n_S$ and $R$ as a function  of $m_{Higgs}$ as well as through their predictions for the consistency relations.\\\\ In order to ensure consistent results, we have studied the SM Higgs inflation variants considering the full dynamics of the background and perturbations of the Higgs field during inflation and compute the RG improved Higgs potential by including the two-loop $\\beta$-functions for the SU(2) $\\times$ SU(1) gauge couplings, the SU(3) strong coupling, the top Yukawa coupling, the Higgs quadratic coupling and the Higgs field anomalous dimension. Assuming that all SM Higgs inflation variants are consistent theories, we use the MCMC technique to derive constraints on the inflationary parameters and the Higgs boson mass from their fit to WMAP7+SN+BAO data set.  We show that the SM Higgs inflation variants lead to significant constraints on $n_s$, $R$ and $m_{Higgs}$.\\\\ From the analysis of the confidence regions in $n_S$-$m_{Higgs}$ plane we conclude that in order to discriminate the SM Higgs inflation variants by using these parameters the measurement by LHC of $m_{Higgs}$ and an improved detection of $n_S$ are required. If PLANCK satellite, designed to measure $n_S$ to an accuracy of $4 \\times 10^{-4}$, will find $n_S$ smaller then 0.97,  while LHC will find a Higgs mass significantly smaller then 151 GeV, the $\\zeta$-inflation and $E$-inflation models would be ruled out. \\\\ We also find a striking difference for the predicted values for $R$ by different models. While for $E$-inflation and $\\zeta$-inflation the predicted values for $R$ are low, the G-inflation model predicts $R$ values enough large to be detected by the PLANCK satellite. We conclude that if PLANCK will detect a value of  $R \\sim 0.1$ while  LHC will find a Higgs mass significantly smaller then 151 GeV, the $\\zeta$-inflation and $E$-inflation models would be ruled out.  \\\\ We also show that the consistency relations are distinct enough to differentiate the SM Higgs inflation variants."
    },
    "1107/1107.4216_arXiv.txt": {
        "abstract": "We report on the MAXI GSC X-ray monitoring of the Crab nebula and pulsar during the GeV gamma-ray flare for the period of 2010 September 18$-$24 (MJD 55457$-$55463) detected by AGILE and Fermi-LAT. There were no significant variations on the pulse phase averaged and pulsed fluxes during the gamma-ray flare on time scales from 0.5 to 5 days. The pulse profile also showed no significant change during this period. The upper limits on the variations of the pulse phase averaged and pulsed fluxes for the period MJD 55457.5$-$55462.5 in the 4$-$10 keV band are derived to be 1 and 19\\%, respectively, at the 90\\% confidence limit of the statistical uncertainty. The lack of variations in the pulsed component over the multi-wavelength range (radio, X-ray, hard X-ray, and gamma-ray) supports not the pulsar but the nebular origin for the gamma-ray flare. ",
        "introduction": "The Crab nebula has been the standard candle in high energy X-ray and gamma-ray astronomy. The flux and spectrum in these energy ranges have been expected to be steady over years. Surprisingly, AGILE and Fermi-LAT reported a flare for the period of 2010 September 18$-$24 in the GeV gamma-ray energy range \\citep{Tavani+2010, Buehler+2010, Tavani+2011, Abdo+2011}. The first half of the flare (September 18$-$21) was detected by both AGILE and Fermi-LAT, while that of the second half (September 21$-$24) was detected only by Fermi-LAT. The half-day binned light curve of Fermi-LAT exhibited three sub-flares during these periods \\citep{Balbo+2011}.  INTEGRAL observed the Crab nebula from 10:32 on September 12 to 12:48 on September 19 (UT) for calibration purposes, which covered the first fifth of the gamma-ray flare. It detected no significant flux increase in the 20$-$400 keV range \\citep{Ferrigno+2010a}. Swift BAT detected no variations over the uncertainty of 5.5\\% at the 1-$\\sigma$ limit in the 15$-$50 keV range for the period of 2010 September 19$-$21 \\citep{Markwardt+2010}. Radio observations of the Crab pulsar showed no evidence of a pulsar glitch and also no change on the pulsed flux as well as the pulse profile \\citep{Espinoza+2010}. ARGO-YBJ reported an excess of events (4 $\\sigma$) from a direction consistent with the Crab nebula, corresponding to a flux about 3$-$4 times higher than usual, at the median energy of about 1 TeV \\citep{Aielli+2010}. On the other hand, MAGIC and VERITAS reported no significant enhancement in the TeV gamma-ray flux during the GeV flare \\citep{Mariotti+2010, Ong+2010}. The Swift XRT follow-up observation of 1-ks exposure starting on September 22 at 16:42 (UT), which corresponds to the third sub-flare recognized in the Fermi-LAT light curve \\citep{Balbo+2011}, showed no changes of flux, spectrum or pulse profile \\citep{Evangelista+2010}. Swift XRT detected no active AGN near the Crab nebula \\citep{Heinke 2010}. The INTEGRAL follow-up observation from 21:05 on September 22 to 11:12 on September 23 (UT), approximately corresponding to the end of the flare, found no significant change in the pulse profile in the 20$-$40 keV band \\citep{Balbo+2011}.  After the gamma-ray flare, the RXTE PCA follow-up observation on September 24 showed no changes of flux, spectrum nor pulse profile \\citep{Shaposhnikov+2010}. The follow-up observation in a near-infrared wavelength on September 24 showed no variation in the Crab pulsar in the J and H bands \\citep{Kanbach+2010}. The follow-up observations by Chandra on September 28 and the HST on October 2 revealed an anvil feature close to the base of the pulsar jet and elongated striation at the distant place \\citep{Tennant+2010, Ferrigno+2010b, Horns+2010, Tavani+2011}. These features are thought to indicate a particle acceleration and the origin of the GeV gamma-ray flare \\citep{Tavani+2011}. Two further GeV gamma-ray flare episodes, detected by AGILE on 2007 October and by Fermi-LAT on 2009 February, were reported \\citep{Tavani+2011, Abdo+2011}. None of these three gamma-ray flares showed any changes of the pulsed component in the GeV energy range.  Another surprising thing is the long-term variability of the Crab nebula in the hard X-ray bands reported by \\citet{Wilson-Hodge+2011}. The total nebula flux was found to decrease from 2008 August to 2010 July and the fractional decline was larger in higher energy ranges. On the other hand, the pulsed flux in the $3.2-35$ keV band of RXTE PCA decreased steadily at $\\sim 0.2$\\% yr$^{-1}$, consistent with the pulsar spin-down, indicating that the observed X-ray variability would originate not from the pulsar but from the nebula.  MAXI has been monitoring the Crab nebula since the beginning of the mission on 2009 August 15 (UT), which covered the whole gamma-ray flare in 2010 September. Here, we report on the MAXI monitoring of the pulse phase averaged and pulsed fluxes of the Crab nebula, as well as the pulse profile. ",
        "conclusions": "\\label{sec: conclusion}  We report on the MAXI GSC observation of the Crab nebula during the GeV gamma-ray flare. We successfully detected the pulsation of the Crab pulsar during the simultaneous period, and conclude that there is no evidence for changes in the pulse profile, pulsed flux and pulse phase averaged flux during the gamma-ray flare. We obtained an upper limit on the variation of the pulse phase averaged flux of $1$\\% at a 90\\% confidence limit of the statistical uncertainty from the best-fit linear function during the 5 day interval of the gamma-ray flare in the 4$-$10 keV band. During the same period in the same energy band, we also obtained an upper limit on the variation of the pulsed flux of $19$\\% at a 90\\% confidence limit of statistical uncertainty from the mean level.  The MAXI GSC simultaneous observation with the gamma-ray flare is uniquely important to constrain the origin of the flare, in contrast to the follow-up observations performed after the cease of the gamma-ray flare. The lack of changes on the pulsed component in the X-ray (this work), as well as those in radio \\citep{Espinoza+2010}, hard X-ray (Super-AGILE observation at the flare on 2007 shown in \\citet{Tavani+2011}) and gamma-ray bands \\citep{Tavani+2011, Abdo+2011}, supports the nebular origin for the gamma-ray flare as proposed in several papers \\citep{Tavani+2011, Abdo+2011, Bednarek Idec 2010, Yuan+2010, Komissarov Lyutikov 2010}. In spite of the large flux increase of factor five in the GeV energy region \\citep{Abdo+2011}, we constrain a limit on the variation of the nebula flux in the X-ray band. This provides valuable information to construct theoretical models for the gamma-ray flare of the Crab nebula.  \\bigskip  We are grateful to the members of the MAXI operation team. We acknowledge the use of the Crab ephemeris provided at the web site of the Jodrell Bank Centre for Astrophysics \\citep{Lyne+1993}. This research was partially supported by the Ministry of Education, Culture, Sports, Science and Technology (MEXT), Grant-in-Aid No.19047001, 20041008, 20540230, 20244015 , 20540237, 21340043, 21740140, 22740120, and Global-COE from MEXT ``The Next Generation of Physics, Spun from Universality and Emergence'' and ``Nanoscience and Quantum Physics.''"
    },
    "1107/1107.1096_arXiv.txt": {
        "abstract": "\\noindent Current studies of the peculiar velocity flow field in the Local Universe are limited by the lack of detection of galaxies behind the Milky Way. The contribution of the largely unknown mass distribution in this ``Zone of Avoidance'' (ZoA) to the dynamics of the Local group remains contraversial. We have undertaken a near infrared (NIR) survey of \\HI~detected galaxies in the ZoA. The photomety derived here will be used in the NIR Tully-Fisher (TF) relation to derive the peculiar velocities of this sample of galaxies in the ZoA. ",
        "introduction": "The mass density field can be inferred from the peculiar velocities of galaxies, independent of any a priori assumption on the relation or bias between visible and dark matter. The determination of the peculiar velocity field requires relatively large and uniform galaxy samples with high-fidelity distance measurements independent of the observed redshift. This can for instance be realized through the Tully-Fisher relation \\citep[][]{TF1977}. However, peculiar velocity surveys are hardly feasible in the ZoA where the obscuring effects of dust and stars in the Milky Way prevent the identification of galaxies across $10-20$\\% of the sky. While this problem can be circumvented by statistical interpolation of the mass distribution adjacent to the ZoA, various papers such as those by \\cite{Kolatt1995} and \\cite{Loeb2008} suggest that these solutions are inadequate and require unknown mass distributions to satisfactorily explain the peculiar motion of the Local Group with respect to the Cosmic Microwave Background. Several such dynamically important structures, including the Great Attractor \\citep[GA;][]{Lynden-Bell1988} and Local Void \\citep{TullyFisher1987} are known to reside within the ZoA.  The 2MASS Tully-Fisher survey \\citep[2MTF;][]{Masters+2008,Masters2008b} is  an ongoing project that aims at measuring TF distances for all bright inclined spirals in the 2MASS Redshift Survey  \\citep[2MRS;][]{Huchra2005}.  The use of the NIR bands will reduce the extinction effects due to the ZoA. Together with the new high-fidelity \\HI~measurements that are being obtained, the 2MTF will provide better statistics for the study of the local peculiar velocity flow field over the whole sky than ever before. However, there still remains a significant part of the sky that will be excluded by the 2MTF. We aim to provide TF data for galaxies in the ZoA which the 2MTF avoids.  These data will be used to measure, for the \\textit{first time}, the peculiar velocity flow field \\textit{within} the southern ZoA. ",
        "conclusions": ""
    },
    "1107/1107.3570_arXiv.txt": {
        "abstract": "It is believed that soon after the Planck time, Einstein's general relativity theory should be corrected to an effective quadratic theory. In this work we present the $3+1$ decomposition for the zero vorticity case for arbitrary spatially homogenous spaces. We specialize for the particular Bianchi $I$ diagonal case. The $3-$ curvature can be understood as a generalized potential, and the Bianchi $I$ case is a limiting case where this potential is negligible to the dynamics. The spirit should be analogous, in some sense to the BKL solution. In this sense, a better understanding of the Bianchi $I$ case could shed some light into the general Bianchi case. ",
        "introduction": "The semi-classical theory consider the back reaction of quantum fields in a classical geometric background. It began about forty years ago with De Witt Ref. \\refcite{DeWitt}, and since then, its consequences and applications are still under research, see for example Ref. \\refcite{Hu}.  Different from the usual Einstein-Hilbert action, the one loop effective gravitational action surmounts to quadratic theories in curvature, see for example Refs. \\refcite{DeWitt,liv}. It is the gravitational version of the Heisenberg-Euler electromagnetism. As it is well known, vacuum polarization introduces non linear corrections into Maxwell electrodynamics,\\cite{schwinger} first obtained by Heisenberg- Euler.\\cite{heisenberg}  This quadratic gravity was previously studied by Starobinsky.\\cite{S} It is of interest for example in the context of the final stages of evaporation of black holes, inflationary theories,\\cite{coule} in the approach to the singularity,\\cite{montani} and also in a more theoretical context.\\cite{cotsakis}  The effective gravity, was apparently first investigated in Tomita's article Ref. \\refcite{berkin} for general Bianchi $I$ spaces. They found that the presence of anisotropy contributes to the formation of the singularity. Berkin's work shows that a quadratic Weyl theory is less stable than a quadratic Riemann scalar $R^2$. In particular, in the very interesting article Ref. \\refcite{barrow-hervik}, Barrow and Hervik addressed the anisotropic cases of Bianchi $I$ and $II$. The most interesting results of Barrow and Hervik are the exact solutions for quadratic theories of the same type investigated in this present article. Instead of the metric, the field equations in Ref. \\refcite{barrow-hervik} are written in a different set of variables which by now is a well known procedure used in the context of dynamical systems in cosmology.\\cite{dsc}. Homogenous solutions in the context of quadratic gravity was also addressed by Ref. \\refcite{resto}. There is an interesting article by Saridakis, Ref. \\refcite{saridakis} in which the Kantowski Sachs anisotropic case is addressed for quadratic gravity. Schmidt does a review of higher order gravity theories in connection to cosmology.\\cite{hjs}  Also in the context of quadratic theories we have the alternative of Gauss-Bonnet type $F(G)$ by Odintsov, Nojiri and collaborators Ref. \\refcite{odintsov06} (for recent reviews see Ref. \\refcite{odintsov}). When the action depends on a arbitrary function of the Gauss-Bonnet term, it is not a top invariant and a consistent dynamic follows from it. Theories of $R^2$ type are also investigated, for example in Ref. \\refcite{shapiro-odintsov}.  In our case the Gauss-Bonnet term is understood as a surface term and discarded.  We have previously addressed the Bianchi $I$ solutions in Refs. \\refcite{sandro,daniel-sandro}. Stability is also discussed there. More recently we have submitted two articles, one in which we investigate how the approach to Minkowski space occurs in the Bianchi $I$ case. \\cite{daniel-marcio-jcarlos} In the other article we present the Bianchi $VII_A$ solutions Ref. \\refcite{juliano-daniel}. In Ref. \\refcite{daniel-marcio-jcarlos} it is obtained that the solution is a superposition of a pure tensorial component, and a pure scalar component. Thus the approach to Minkowski space involves the production of scalar and tensorial gravitational waves. Speculations about the validity of the semi-classical argument can be raised, anyway, in a certain sense the quadratic counterterms which we consider are the most natural ones expected from the renormalization of a quantum field.\\cite{christensen}  The purpose of this article is to cast the dynamical equations of motion in a 3+1 decomposition for the zero vorticity case. A time like and geodesic vector can be defined and we impose the homogeneity on the 3-space. Apparently it is the first time the equations are presented in this form. The Bianchi $I$ case is revisited but now the equations are written in an analytical fashion. We intend to use the expressions obtained in this work in our future works. Most of the article is there just for completeness.  The article is organized as follows. In section 2 a brief exposition of the vacuum polarization by the external gravitational field. In section 3 we present the $3+1$ decomposition for any Bianchi type. In the section 4 we specialize to the Bianchi I case. And the conclusions are presented in section 5. ",
        "conclusions": "The semi-classical theory consider the back reaction of quantum fields in a classical geometric background. It began about forty years ago with De Witt Ref. \\refcite{DeWitt}, and since then, its consequences and applications are still under research, see for example Ref. \\refcite{Hu}.  Different from the usual Einstein-Hilbert action, the one loop effective gravitational action surmounts to quadratic theories in curvature, see for example Refs. \\refcite{DeWitt,liv}. It is the gravitational version of the Heisenberg-Euler electromagnetism. As it is well known, vacuum polarization introduces non linear corrections into Maxwell electrodynamics,\\cite{schwinger} first obtained by Heisenberg- Euler.\\cite{heisenberg}  In this article we present the decomposition $3+1$ for arbitrary spatially homogenous space times for this particular effective quadratic theory. It turns out that for this particular gravitational theory \\eqref{acao}, every Einstein space satisfying $R_{ab}=g_{ab}\\Lambda/2$ is an exact solution, this includes of course the vacuum case $\\Lambda=0$. Anyway the most interesting is the BKL oscillatory approach to the singularity which occurs for vacuum in Einstein's context.\\cite{bkl} Since the mixmaster solution is a vacuum $R_{ab}=0$ solution for the Bianchi $IX$ case it will occur exactly the same in the effective theory considered in \\eqref{acao}.  In this sense a better understanding of the Bianchi $I$ case for the quadratic theories seems attractive. The $3$ curvature also can be understood as a potential and while approaching the singularity its influence in the dynamics can decrease almost to zero. Thus, an arbitrary Bianchi solution of the quadratic gravity can approach very much the Bianchi $I$ case. We intend to use this $3+1$ decomposition, numerically, in future works. Particular Bianchi $I$ analytic exact solutions for a quadratic theory identical to this one were already found by Ref. \\refcite{barrow-hervik}. Also for $f(R)$ gravity models including the $R+R^2$ case, Bianchi I space-times were studied in Ref. \\refcite{gurovich}, where it was shown that equations for the anisotropic part of the metric can be integrated.  We emphasize that the approach of this present work is facing the theory \\eqref{acao}  as an effective, and classical theory. While considering it as a candidate for a quantized gravity does introduces ghosts. Which together with tachyons and the additional degrees of freedom in contrast to Einstein's theory, makes it improbable that the BKL solution will be the generic one in this case."
    },
    "1107/1107.5740_arXiv.txt": {
        "abstract": "We examine sources of scatter in scaling relations between galaxy cluster mass and thermal Sunyaev-Zeldovich (SZ) effect using cluster samples extracted from cosmological hydrodynamical simulations. Overall, the scatter of the mass-SZ scaling relation is well correlated with the scatter in the mass-halo concentration relation with more concentrated halos having stronger integrated SZ signals at fixed mass. Additional sources of intrinsic scatter are projection effects from correlated structures, which cause the distribution of scatter to deviate from log-normality and skew it towards higher inferred masses, and the dynamical state of clusters. We study the evolution of merging clusters based on simulations of 39 clusters and their cosmological environment with high time resolution. This sample enables us to study for the first time the detailed evolution of merging clusters around the scaling relation for a cosmologically representative distribution of merger parameters. Major mergers cause an asymmetric scatter such that the inferred mass of merging systems is biased low. We find mergers to be the dominant source of bias towards low inferred masses: over 50\\% of outliers on this side of the scaling relation underwent a major merger within the last Gigayear. As the fraction of dynamically disturbed clusters increases with redshift, our analysis indicates that mergers cause a redshift-dependent bias in scaling relations. Furthermore, we find the SZ morphology of massive clusters to be well correlated with the clusters' dynamical state, suggesting that morphology may be used to constrain merger fractions and identify merger-induced outliers of the scaling relation. ",
        "introduction": "Clusters of galaxies are the most massive gravitationally bound objects in the universe, which makes them an important tool for cosmology: among other tests, their abundance provides information on the gravitational growth of structures and is regulated by the initial density field, gravity, and the expansion history of the universe, which critically depend on the underlying cosmology. Thus number counts of clusters, for which masses and redshifts are known, can be used to constrain cosmological parameters \\citep[see][for a recent review]{Allen11}.\\\\ To relate observed number counts to theoretical predictions of the cluster mass function, these experiments need to infer cluster masses from observables. The thermal Sunyaev Zeldovich (SZ) effect, the signature of inverse Compton scattering of cosmic microwave background photons with hot cluster electron, is thought to provide an excellent mass proxy as the SZ signal is proportional to the total thermal energy of a cluster and is thus less affected by physical processes in the cluster core which can largely affect the X-ray luminosity. This is confirmed by simulations \\citep[e.g.][]{Nagai06, Shaw08, Battaglia10, Sehgal10} finding the scatter in the mass - SZ scaling relation to be of order 5 - 10\\%. Furthermore, the SZ effect is not subject to surface brightness dimming and has a very weak redshift dependence, making it an ideal probe to study galaxy clusters at high redshift.\\\\  Currently several large surveys are starting to detect hundreds of galaxy clusters through their SZ signal \\citep{Vanderlinde10,Marriage10, Planck11} and derive cosmological constraint based on these samples \\citep{Andersson10, Sehgal10b, Williamson11}. To exploit the statistical power of these upcoming cluster samples, the mapping between SZ signal and cluster mass needs to be well understood. Observations find normalization and slope of the scaling relations between SZ signal and lensing derived masses \\citep{Marrone11}, or between SZ signal and X-ray properties \\citep{Planck11b, Planck11c} to be consistent with self-similar scaling and predictions from simulations.\\\\  Due to the steep slope of the cluster mass function, competitive cosmological constraints from these experiments require information about the distribution and redshift evolution of scatter in the mass scaling relation \\citep[e.g.][]{Majumdar04, Lima05,Shaw10}. As the true cluster mass and other physical cluster properties which may bias the mass proxy are unobservable, and as the noise and biases in the different mass estimators may be correlated, characterizing the intrinsic scatter in any of these scaling relation is difficult to obtain from observations. Hence the sources and distribution of scatter in different mass estimators are mainly studied through simulations and mock observations \\citep[e.g.][]{Rasia06,Nagai07,Shaw08,Becker10,Yang10,Fabjan11}.\\\\  In this work we focus on the effect of merging events on the SZ signal of a galaxy cluster. As clusters form through merging of smaller objects these are frequent and disruptive events, which may alter the physical state of the involved clusters significantly. Hence merging clusters may deviate from the scaling relations observed in relaxed clusters and, as the fraction of morphologically disturbed clusters increases with redshift, cause a redshift dependent scatter or bias in the mass scaling relation. Simulations of binary cluster mergers \\citep{Randall02,Poole06, Poole07, Wik08} find that the X-ray luminosities, temperatures, SZ central Compton parameters and integrated SZ fluxes increase rapidly during the first and second passage of the merging clusters. The clusters temporarily drift away from mass scaling relations and return to their initial scaling relation as the merging system virializes. These transient merger boosts found in binary mergers and some observations \\citep{Smith03} can scatter the inferred masses towards higher values and thus bias the derived cosmology towards a higher normalization of the power spectrum, $\\sigma_8$, and lower matter density \\citep{Randall02,Smith03,Wik08,Angrick11}. On the other hand, mergers increase the non-thermal pressure support \\citep{Rasia06, Lau09, Battaglia10} found in cluster outskirts, and due to partial virialization merging clusters can appear cooler than relaxed clusters of the same mass \\citep[e.g.][]{Mathiesen01}. For a cluster sample extracted from cosmological simulations, \\citet{Kravtsov06} find the X-ray temperatures of morphologically disturbed clusters to be biased, while the X-ray derived SZ-equivalent $Y_X$ shows no significant correlation with cluster structure. Comparing X-ray and SZ to weak lensing derived masses, \\citet{Okabe10} and \\citet{Marrone11} found undisturbed clusters to have of order $\\sim\\!40\\%$ higher weak lensing masses than disturbed clusters at fixed $T$ and $Y_{\\mr{SZ}}$, and $\\sim\\!20\\%$ higher weak lensing masses at fixed $Y_{\\rm X}$.\\\\  Our goal is to isolate how mergers in a cosmological context affect the SZ signal of clusters, and if merging cluster can be detected as outliers of scaling relations. This extends previous work, as our analysis includes both multiple mergers with realistic distributions of orbits and mass ratios, and full SPH treatment of gas physics with radiative cooling, star formation and supernova feedback. The simulations and the cluster sample are described in Sect. \\ref{sec:simul}. We discussion the best-fit scaling relations and their scatter in Sect. \\ref{sec:scaling}. The effect of merging events of the clusters SZ signal is quantified and the evolution of merging clusters with respect to the scaling relations is discussed in Sect. \\ref{sec:merger}. In Sect. \\ref{sec:morph} we investigate if the dynamical state of clusters can be inferred from the morphology of the SZ signal. We summarize our results and conclude in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} Using projected Compton $y$ maps of galaxy clusters extracted from cosmological hydrodynamical simulations, we analyze the clusters' thermal SZ signal and its scaling relation with cluster mass. We study the detailed time evolution of a sample of 39 clusters around the scaling relation using simulations with outputs closely spaced in time. Compared to previous studies, which focused either on the evolution of isolated, idealized mergers or on large samples of clusters at widely spaced redshifts, this sample enables us to isolate the effect of merging events for a cosmologically representative distribution of merger orbits, mass ratios, and impact parameters. Our main results can be summarized as follows: \\begin{enumerate} \\item The best fit scaling relations to the integrated $Y_{200}$ signal of these clusters are close to self-similar predictions and agree well with other simulations that include comparable gas physics. \\item The scatter around these scaling relations is small (of order 10\\% scatter in mass at fixed $Y_{200}$) and it is overall well correlated with the scatter in halo concentration, such that more concentrated halos have larger $Y$ signal at fixed mass. \\item The scatter in the scaling relation deviates from a log normal distribution and is skewed towards clusters with $Y$ signals larger than expected from their mass. We find projection effects due to nearby structures to be an important source of this upward scatter. However, due to the small magnitude of the scatter in the mass scaling projection effects are not expected to be a significant contamination for cosmological constraints from SZ cluster surveys. \\item Merging clusters fall below the scaling relation, such that their inferred masses are biased low. More quantitatively, we find that within a Gyr following a merger, clusters are twice as likely as the average cluster to undergo a phase during which their inferred mass is biased low by more than 10\\%. \\item We identify merging events to be a major source of downward scatter in the scaling relation: a large fraction of clusters whose inferred masses are biased low recently underwent a merger (c.f. Fig.~\\ref{fig:lookfw}). \\item For massive clusters, we find the morphology of SZ maps to be well correlated with deviations from the scaling relation. While the robustness of this result with respect to noise and imaging artifacts requires further analysis, it suggests that SZ morphology may be useful to reduce the scatter of mass estimates, and to infer merger rates of massive halos and hence test theories of halo formation. \\end{enumerate} Our analysis of the time evolution of merging events is in agreement with the conclusions drawn from earlier studies comparing morphologically disturbed and undisturbed clusters in cosmological simulations at fixed redshifts \\citep[e.g.][]{Mathiesen01, Kravtsov06,Nagai06,Jeltema08, Ventimiglia08}. Specifically, it supports the hypothesis that for a cosmological distribution of merger parameters partial virialization and non-thermal pressure support due to mergers are more important than merger boosts found in simulations of direct collisions between mergers. For simulated clusters the intrinsic scatter in the scaling relation and the mass segregation between morphologically relaxed and disturbed clusters are significantly smaller than recent observational results based on SZ measurements, X-ray morphology and weak lensing inferred masses \\citep{Marrone11}. However, as these authors note, the observed scatter is in agreement with the scatter expected in weak lensing mass measurements \\citep{Becker10}. Similarly, the mass segregation is enhanced by the sensitivity of weak lensing mass estimates to cluster triaxiality, and these observational constraints on the intrinsic scatter and bias in SZ mass estimates are limited by the accuracy of weak lensing mass reconstruction.\\\\  Further complications arise when inferring cluster masses from SZ observations as most $Y$ measurements are derived from fitting parametric profiles \\citep[e.g.][]{Nagai07p,Arnaud10} to the data which assume radial symmetry \\citep[but see][for alternate methods and discussions]{Plagge10,Marrone11,Sayers11}.The distorted geometry of merging clusters may introduce additional scatter to mass estimates derived from profile fits, but an experiment specific analysis of such effects is beyond the scope of this work. \\\\  An additional limitation of our analysis is the range of non-gravitational physics included in the simulations. While recent studies show the impact of AGN-feedback on overall cluster profiles and scaling relations \\citep{Sijacki07,Puchwein08,Battaglia10, Fabjan11}, this mainly affects the cluster center. Consequently, we do not expect AGN feedback to significantly alter the slow virialization of newly accreted material at larger radii, which we found to be the main source of scatter during merging events. In the cluster outskirts, electrons and ions are not in thermal equilibrium. \\citet{Rudd09} and \\citet{Wong09} show that detailed treatment of the multi-temperature structure of the intracluster medium leads to a significant suppression of electron temperature and SZ signal. Based on a sample of three simulated cluster, \\citet{Rudd09} find this effect to be especially pronounced in clusters undergoing major mergers. Under specific conditions, this effect may cause a bias of up to $5\\%$ in integrated $Y$, corresponding to an additional negative bias of about $ 3\\%$ in the inferred mass of merging clusters.\\\\  Overall, we find that merger events cause a temporary negative bias in inferred cluster mass of order $10\\%-15\\%$. Due to the increased fraction of recently merged objects at higher redshift, we conclude that this merger bias should be accounted for when modeling the redshift evolution in the scatter of scaling relations."
    },
    "1107/1107.2356_arXiv.txt": {
        "abstract": "We propose a new way to detect individual bright Ultra-High Energy Cosmic Ray (UHECR) sources above background if the Galactic Magnetic Field (GMF) gives the main contribution to UHECR deflections~\\cite{SearchForSources}. This method can be directly applied to maps given by experiments. It consists in starting from at least two high energy events above 6\\,$\\times$\\,10$^{19}$\\,eV, and looking at lower energy tails. We test the efficiency of the method and investigate its dependence on different parameters. In case of detection, the source position and the local GMF deflection power are reconstructed. Both reconstructions are strongly affected by the turbulent GMF. With the parameters adopted in this study, for 68\\% of reconstructed sources, the angular position is less than one degree from the real one. For typical turbulent field strengths of 4\\,$\\mu$G at the Earth position and 1.5\\,kpc extension in the halo, one can reconstruct the deflection power with 25\\% precision in 68\\% of cases. ",
        "introduction": "The HiRes experiment first observed the cutoff in the spectrum~\\cite{spectrum_HiRes}. The UHECR sources are still unknown. Some possible candidates can be studied from the point of view of acceleration mechanisms -see for example Refs.~\\cite{Hillas:1985is,Ptitsyna:2008zs,Neronov:2007mh}.\\\\ Their detection is difficult due to the UHECR deflections in the extragalactic and Galactic magnetic fields. One can either wait for enough data at the highest energies to find a class of UHECR sources, or search for the first individual brightest sources. In this study, we investigate this second possibility.\\\\ We propose here a method to find Ultra-High Energy (UHE) proton or light nuclei sources on top of background and reconstruct their positions. We show that one has to start from at least two events with energies above $10^{19.8}$\\,eV\\,$\\sim 6 \\times 10^{19}$\\,eV before looking at lower energy tails, in order to avoid confusion by the background. We investigate the performance of this method depending on the different relevant parameters.\\\\ We assume as in~\\cite{Dolag:2004kp} that UHECR deflections due to the extragalactic magnetic fields are negligible compared to those due to the GMF. For an example of non negligible extragalactic contributions, see~\\cite{Sigl:2004yk}. The GMF is divided in two components: regular and turbulent. For the regular one, we take the Prouza and Smida model~\\cite{PS,GMF}. It consists of a thick disk, and of toroidal and dipolar fields. Following the latest knowledge on the GMF~\\cite{Waelkens:2008gp,Sun:2007mx}, we modify some parameters and take an exponentially decaying profile along the Galactocentric radius. The field strength close to the Sun is $B_{\\rm reg}=1-3\\,\\mu$G (currently admitted strength: $B_{\\rm reg} \\simeq 2\\,\\mu$G). The model for the turbulent component is described in~\\cite{SearchForSources}. The root mean square (RMS) value of the turbulent field decays as~$1/r$ along the Galactocentric radius~$r$, and exponentially along the direction orthogonal to the Galactic plane. The plane thickness is set here to 1.5\\,kpc. The turbulent field strength close to the Sun is $B_{\\rm turb}=2-8\\,\\mu$G (currently admitted strength: $B_{\\rm turb} \\simeq 4\\,\\mu$G).\\\\  \\begin{figure}[!t] \\centering \\includegraphics[width=0.5\\textwidth]{icrc0422_fig01.eps} \\caption{Projections of the arrival directions of cosmic rays (CR) emitted by a source (cross) in the plane tangent to the celestial sphere, and centered on the highest energy event (filled square in (0,0)) emitted by the source. Deflections X and Y measured on two orthogonal axes, in degrees. Open symbols for source CR deflected by the regular GMF only. Filled symbols for the same CR when the turbulent component is also taken into account. Shapes correspond to CR energies (triangles: between $10^{19.0}$ and $10^{19.3}$\\,eV, circles: between $10^{19.3}$ and $10^{19.6}$\\,eV, and squares: above $10^{19.6}$\\,eV).} \\label{TurbulentComponent} \\end{figure}  For source events, we assume power law acceleration spectra, with a power law index equal to 2.2 and a maximum energy set to 10$^{21}$\\,eV. In most cases, we take a proton source and set its distance to the Earth to $50-100$\\,Mpc. Protons are propagated to the observer by using the results of Ref.~\\cite{Kachelriess:2007bp} for energy losses, which creates a ``bump'' in the source spectrum~\\cite{bump}. Then, the generated cosmic rays (CR) are deflected in the regular GMF (see e.g.~\\cite{Harari2}), and in the turbulent GMF. For more details on turbulent GMF deflections, see Refs.~\\cite{TurbulentComponent1,SearchForSources}. The effects of both components are shown in Fig.~\\ref{TurbulentComponent}: The regular component deflects the CR along a curve and the turbulent one spreads them ``randomly'' around it in a sector-shaped region. When there are many events with the same energies, they are not spread uniformly, as discussed in References~\\cite{Harari2,TurbulentComponent1,Golup:2009zg}.\\\\ Results depend on many parameters of the model and on the considered location on the sky. However, in practice, there are only two essential parameters: The local deflection powers -see Ref.~\\cite{SearchForSources}- of the regular GMF (denoted $D$) and of the turbulent GMF.\\\\ Background events are simulated according to the energy spectrum measured by HiRes~\\cite{spectrum_HiRes} and the exposure of Telescope Array~\\cite{exposure}. ",
        "conclusions": "We have proposed and studied here an alternative way to detect individual bright UHECR sources, if the UHECR deflections due to the turbulent GMF are not too large compared to the deflections due to the regular component. We generated two sets of events: Events from a source with a $1/E^{2.2}$ injection spectrum that we propagated from the source to the Earth, and events from some background generated according to the HiRes spectrum and the Telescope Array exposure. We mixed both sets of events, tried to detect the source and checked how often we got confused by detecting some background.\\\\ We showed that one should look for doublets of events with $E>10^{19.8}$\\,eV, when using the depicted method.\\\\ We studied in section~\\ref{sectiondependance} the dependence on several parameters: Unknown parameters from the magnetic fields, the source, the sky anisotropy at high energies and the experimental statistics. We showed in section~\\ref{sectionreconstruction} that with the parameters and assumptions adopted here, the precision on the source reconstruction would be dominated by the current experimental angular resolution in 68\\% of cases. The local deflection power can be known up to 25\\% in 68\\% of cases.\\\\ In the future, this method can be applied to experimental data so as to find bright UHECR proton or light nuclei sources."
    },
    "1107/1107.0353_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sect:intr}  Astronomy is a science driven by discovery\\footnote{http://www.gmto.org/sciencecase/GMT-ID-01404-GMT\\_Science\\_Case.pdf}, and the essential components in the astronomical discoveries are telescopes and instruments. Since 1990, the world's largest optical telescopes with 8 to 10 m in diameter have appeared and gave birth to new innovations in astronomy. After 10 -- 20 years of use by the largest optical telescopes, the necessity for larger telescope has increased and we are witnessing the development of 25 -- 42 m extremely large telescopes (ELTs).  Table \\ref{tab1} lists the current largest, ground-based optical telescopes (diameter D $\\ge$8 m). There are five, more active telescopes mainly built in the 20th century : USA's two 10 m Keck telescopes, four 8.2 m Very Large Telescope (VLT) of the European Southern Observatory (ESO), two 8.1 m Gemini telescopes of the consortium of USA, UK, Canada, Chile, Australia, Brazil, and Argentina, 8.2 m Subaru telescope of Japan, and 9.2 m Hobby-Eberly Telescope (HET). The other three telescopes at the lower part of Table \\ref{tab1} are built later in the 21st century, and are in the early part of their operations, which are the 10.4 m Gran Telescopio Canarias (GTC), the 10 m South African Large Telescope (SALT), and the two 8.4 m (on a single mount) Large Binocular Telescope (LBT).   \\begin{table*} % \\scriptsize{ \\begin{center} \\caption{Current Ground-based Large Optical Telescopes (diameter $\\ge$8 m)}\\label{tab1} \\begin{tabular}{cccc} \\hline\\hline {Telescope} & {Diameter} & {Partner Country/Institute} & {Operation} \\\\ &&& {Start Year} \\\\ \\hline Keck & 10 m (36 $\\times$ 1.8 m segments) $\\times$ 2 & USA (Caltech, University of California) & 1993, 1996 \\\\ VLT$^{\\rm a}$ & 8.2 m $\\times$ 4 & ESO members (Austria, Belgium, Czech Republic, & 1998, 1999, \\\\ & & Denmark, Finland, France, Germany, Italy, Netherlands, & 2000, 2000 \\\\ & & Portugal, Spain, Sweden, Switzerland, United Kingdom) & \\\\ Gemini & 8.1 m $\\times$  2 (N, S hemispheres) & USA (48\\%), UK (24\\%), Canada (14\\%), Chile (5\\%), & 1999, 2000 \\\\ & & Australia (5\\%), Brazil (2\\%), Argentina (2\\%) & \\\\ Subaru & 8.2 m & Japan & 1999 \\\\ HET$^{\\rm b}$ & 9.2 m (91 segments) & USA (90\\%), Germany (10\\%) & 1999 \\\\ \\hline GTC$^{\\rm c}$ & 10.4 m (36 segments) & Spain (90\\%), University of Florida (5\\%), Mexico (5\\%) & 2008 \\\\ SALT$^{\\rm d}$ & 10 m (91 segments) & Republic of South Africa (35\\%), Poland, HET, USA, & \\\\ & & Germany, New Zealand, UK, India & 2005 \\\\ LBT$^{\\rm e}$ & 8.4 m $\\times$ 2 (single mount) & USA (50\\%), Italy (25\\%), Germany (25\\%) & 2008 \\\\ \\hline \\end{tabular} \\medskip\\\\ \\end{center} $^{\\rm a}$ Very Large Telescope \\\\ $^{\\rm b}$ Hobby-Eberly Telescope \\\\ $^{\\rm c}$ Gran Telescopio Canarias \\\\ $^{\\rm d}$ South African Large Telescope \\\\ $^{\\rm e}$ Large Binocular Telescope \\\\ } % \\end{table*}   Table \\ref{tab2} lists the three planned ELT projects, of which telescope diameters are larger than 20 m. When these state-of-the-art telescopes are completed, how much impact will they bring to the astronomical research? One of the methods to answer this question is to scrutinize the impacts the current largest telescopes have brought since their completions. The analysis of the major scientific papers published by using these telescopes would be the essential part in this area, which will be shown in this paper.   \\begin{table*} % {\\footnotesize % \\begin{center} \\caption{Extremely Large Telescope (ELT) Projects in Progress\\label{tab2}} \\begin{tabular}{ccc} \\hline\\hline {Telescope} & {Diameter} & {Partner Country/Institute} \\\\ \\hline Giant Magellan Telescope (GMT) & 25 m (7 $\\times$ 8.4 m segments) & USA, Korea, Australia \\\\ Thirty Meter Telescope (TMT) & 30 m (492 $\\times$ 1.4 m segments) & USA, Canada$^{\\rm a}$ \\\\ European Extremely Large Telescope (E-ELT) & 42 m (984 $\\times$ 1.4 m segments) & ESO \\\\ \\hline \\end{tabular} \\medskip\\\\ \\end{center} $^{\\rm a}$Japan (Collaborating Institution), China and India (Observers) \\\\ } % \\end{table*}   Counting the number of published papers or the number of citations of specific papers which used certain telescopes is one of the ways to measure the impact and importance of these telescopes or facilities (e.g., \\citet{davoust87, leverington96, schulman97, abt98, abt00, trimble08, crabtree08, trimble09, crabtree11}), and analyzing citations is a method to measure the amount of impact of a certain paper \\citep{stanek08}.  From analysis of 11,831 papers published in 20 journals of astronomy and astrophysics from 2001 to 2003, \\citet{trimble08} suggested that the {\\it Hubble Space Telescope (HST)} is responsible for the largest number of optical papers, while the most frequently cited optical papers come from the Sloan Digital Sky Survey (SDSS), Keck, and the Anglo-Australian Telescope (AAT). \\citet{grothkopf05} also showed that the {\\it HST} surpasses both VLT and Keck in the total number of papers, as well as in the numbers of papers per year (their figure 1; see also \\citet{ringwald03, meylan04, apai10}, but see also \\citet{leverington97b}).  By analysing papers resulting from optical telescopes larger than 2 m diameter published in 1990-1991 and cited in 1993, \\citet{trimble95} and \\citet{trimble96} found that the largest numbers of papers and citations came from 4 m class telescopes, the Canada-France-Hawaii Telescope (CFHT) and AAT, followed by the Cerro-Tololo Inter-American Observatory (CTIO) 4 m, while the largest impact factors (five or more citations per paper per year) came from the University of Hawaii's 2.2 m and the Multi-Mirror Telescope % in Arizona \\citep{trimble08}. The analysis of papers published in 2001 and 2002, which is after the completions of 8 m class telescopes, \\citet{trimble07} showed that the largest optical telescopes are responsible for the largest numbers of papers, while 4 m class telescopes displayed continued fading, except for the infrared United Kingdom InfraRed Telescope (UKIRT) and InfraRed Telescope Facility (IRTF).  From the analysis of 1000 most highly-cited papers published between 1991 and 1998 (125 from each year) and 452 astronomy papers published in {\\it Nature} during $1989-1998$, \\citet{benn01} showed that the bigger the telescope, the more the paper cited, with citation fraction $\\propto {\\rm diameter}^2$. \\citet{trimble05} also suggested that big telescopes produce more papers and more citations per paper than small ones, from the analysis of 2100 papers produced in 2001. \\citet{ahn08} suggested that the amount of papers produced by a large ($D \\sim 3.6 - 10$ m) telescope is roughly proportional to the diameter of its primary mirror (see also \\citet{leverington97a}). They also estimated the numbers of refereed and {\\it Nature/Science} papers that might be produced by the Giant Magellan Telescope (GMT) annually to be 330 and 17, respectively: the former by using the rough equation of $N/D \\sim 14 $ ($N$ is number of refereed papers and $D$ is the diameter of a telescope in meter) and the filled aperture 21.4 m of the GMT, and the latter by using another rough equation of $<n/A> \\sim 0.05$ ($n$ is number of {\\it Nature/Science} papers published by Keck I and each VLT telescope, and $A$ is the collecting area of the primary mirror). \\citet{frogel10} initiated a series of papers to investigate what effects the new facilities, data archives, and means of information change had on astronomical publications, first by analysing the 100 most cited papers in each year from 2000 to 2009.  In this paper, we present an analysis of the publications based on results from the largest ground-based optical telescopes of Keck, VLT, Gemini, Subaru, and HET telescopes during the years of $2000-2009$. Using the data, we try to find (i) the temporal trend of the publications from the above telescopes, and (ii) if there is any correlation in the refereed paper publications and the {\\it Nature/Science} paper publications, where the latter is assumed to be the paragon of high-impact journals. The paper is organized as follows: Sect. \\ref{sect:data} describes the data utilized in this work. Sect. \\ref{sect:results} presents the analysis results : Sect. \\ref{sect:number} focuses on the total number of papers and Sect. \\ref{sect:natsci} focuses on the {\\it Nature} and {\\it Science} papers. Finally, Sect. \\ref{sect:sum} summarizes and discusses the results. ",
        "conclusions": "\\label{sect:sum}  We have analysed the ten year ($2000-2009$) publication record of the current largest (D $>8$ m) ground-based optical telescopes of Keck, VLT, Gemini, Subaru, and HET. During the ten year period, the telescopes of Keck, VLT, Gemini, and Subaru showed increasing numbers of refereed papers and this tendency is still preserved when we divided the number of papers by the number of telescope components (2 for Keck and Gemini, and 4 for the VLT telescopes). Each telescope of the Keck, VLT, Gemini, Subaru, and HET observatories produced 135, 109, 93, 107, and 5 refereed papers, respectively, in 2009. For the ten year period, the number of papers produced by each of the telescopes is largest for the Keck, while the largest slope in the change of the annual number of papers is for the VLT. It is worthwhile to note that the impact of papers based on archival data can have a significant impact on a telescope's productivity. For example, almost half of papers published using $HST$ data are based on at least some archival data\\footnote{http://archive.stsci.edu/hst/bibliography/pubstat.html}, and this could be also a factor for VLT, Gemini, and Subaru as mentioned in \\S 3.1. While the astronomical literature continues to grow exponentially by $2-3$\\% \\citep{frogel10}, 4\\% \\citep{abt98}, 5\\% \\citep{white07,trimble08}, $6-7$\\% \\citep{abt10}, or 8.8\\% \\citep{abt95} annually, we will need more data for the next several years to see if the number of papers produced by using the telescopes will still increase or not, since some of the telescopes (Keck and VLT) show somewhat less publications in 2008 and 2009 than the year of 2007 (see Figure \\ref{fig2} (a)).  For the papers published in the two multi-disciplinary, high-impact journals of {\\it Nature} and {\\it Science} \\citep{frogel10}, we have shown that each telescope of the Keck, VLT, Gemini, and Subaru observatories is producing $2.1 \\pm 0.9$ {\\it Nature} and {\\it Science} papers annually and the rate of these papers among all the refereed papers produced by using that telescope is $1.7 \\pm 0.8$ \\%. Extending this relation obtained from the current largest ground-based optical telescopes, we may be able to conclude that this ratio of the number of {\\it Nature} and {\\it Science} papers over the number of whole refereed papers that will be produced by future ELTs of GMT, TMT and E-ELT will remain similar. If, therefore, one of the future larger telescopes produces, e.g., 330 refereed papers annually, the above $simple$ calculation suggests that $\\sim 6$ {\\it Nature} and/or {\\it Science} papers might be included in these publications annually.  From the comparison of the publication trends of the telescopes, we may conclude the followings : \\begin{itemize} \\item {\\em} While the telescope productivity means papers per telescope, it is expected that the more telescopes of the same kind at the same location, the more synergies occur. This includes the effectiveness of maintenance, less number of observatory personnel, less cost for the facilities and more chances to use the instruments that are made to be attached to the same telescope. Although this fact might not be the critical factor for telescope productivity, the specific example of the VLT is worth noting. The four VLT telescopes currently have the largest number of instruments (12 ; see \\S 3.1, four of the instruments can be used at the same time), largest number of papers among the telescopes considered in this study, and largest slope value ($12.1 \\pm 1.2$) of the fitting of the number of refereed papers over year.  \\item {\\em} The important factors that influence the growth rate of paper production are ramp-up in efficient operations, reliable instruments, useful instruments, and the number of good instruments available at the telescope. The latter point might be supported by the fact that the order of the number of instruments is almost the same as that of the slope values of the fitting of the number of refereed papers over year : VLT    (12, $12.1 \\pm 1.2$), Gemini (11, $9.9 \\pm 0.7$), Keck   (9, $7.7 \\pm 1.2$), Subaru (8, $9.6 \\pm 0.8$), and HET    (3, almost flat), where (number of instruments, the slope value). \\end{itemize}  Although it might not be able to have more than one telescope at a site to maximize the productivity, it is a natural and necessary way to increase the number of good instruments available and to afford the archive to maximize the use of the data. There are also many other items that affect the productivities of telescopes, such as : \\begin{itemize} \\item {\\em} user base of the telescope, \\item {\\em} publication traditions of journals (e.g., US vs European journals) \\citep{schulman97, abt98, abt10, frogel10}, and \\item {\\em} support on the telescope users (pool of the management personnel), \\end{itemize} of which investigation in the future might give us more lessons, although it is not easy to get some of the data (e.g. budget). The first item above, the user base of the telescope, might be correlated with the telescope subscription rate. \\citet{frogel10} showed that (in his figure 1) the membership numbers for the American Astronomical Society (AAS) have stayed flat for the last $10-20$ years (cf. \\citet{abt00}), while those for the International Astronomical Union increased about 20\\% over the same period. This almost constant number of AAS membership shows no correlation with the increase of the numbers of papers of US telescopes of Keck, HET (90\\% portion for US) and Gemini (48\\% portion for US) as shown in this study. This indicates that the analysis of the user base of the large optical telescopes might need the database of actual telescope users, needless to say that of optical astronomers of the countries that operate the telescopes, specifically for VLT, Gemini, and HET which are being operated by two or more countries. The rather detailed analysis on the journal of AJ by \\citet{bracher99} might represent the studies on the publication traditions of journals."
    },
    "1107/1107.0679_arXiv.txt": {
        "abstract": "We have observed the radio nebula surrounding the Galactic LBV candidate G79.29+0.46 with the EVLA at 6~cm. These new radio observations allow a morphological comparison between the radio emission, which traces the ionized gas component, and the mid-IR  emission, a tracer of the dust component.  The IRAC ($8 \\, \\mu m$) and MIPS ($24\\, \\mu m$ and $70\\, \\mu m$)  images have been reprocessed and compared with the EVLA map. We confirm the presence of a second shell at 24$ \\mu m$ and also provide evidence for its detection at $70\\, \\mu m$. The differences between the spatial morphology of the radio and mid-IR maps indicate the existence of two dust populations, the cooler one emitting mostly at longer wavelengths. Analysis of the two dusty, nested shells have provided us with an estimate of the characteristic timescales for shell ejection, providing  important constraints for stellar evolutionary models.  Finer details of the ionized gas distribution can be appreciated thanks to the improved quality of the new 6~cm image,  most notably the highly structured texture of the nebula.  Evidence of interaction between the nebula and the surrounding interstellar medium can be seen in the radio map, including  brighter features that delineate  regions where the shell structure is locally modified. In particular, the brighter filaments in the south-west region appear to frame the shocked southwestern clump reported from CO observations. ",
        "introduction": "Massive stars play a fundamental role in the evolution of galaxies. They are major contributors to the interstellar UV radiation and, via their strong stellar winds, provide enrichment of processed material (gas and dust) and mechanical energy to the interstellar medium. Despite their importance, the details of post-MS evolution of massive stars are still poorly understood. Recent evolutionary models suggest that Luminous Blue Variables (LBVs)  and related transition objects may play a key role in the massive star evolution, representing a crucial phase during which a star loses most of its H envelope  \\citep{lamers01}. More recently, it has also been pointed out that LBVs might be direct Supernovae progenitors \\citep{Vink_06, Smith_08}, enhancing their importance in the framework of stellar evolution  LBVs are luminous (intrinsically bright, $ L \\sim 10^{6} \\,{\\rm L_{\\odot}}$) objects, exhibiting different kinds of photometric and spectroscopic variabilities. They are massive ($M \\sim 20-120 \\,{\\rm M_{\\odot}}$), characterized by  intense mass-loss rates ($10^{-6}-10^{-4}\\,{\\rm M_{\\odot} yr^{-1}}$), which can also occur in the form of eruptive events. Although eruptive events have been witnessed very rarely (e.g., $\\eta$ Car and P Cyg) the presence of extended, dusty  circumstellar nebulae around LBVs  (LBVNe) suggests that they are a common aspect of LBV behavior \\citep{Weis08}. There are, however, many aspects of LBV phenomenon that are not completely understood. Among these are the total mass lost during the LBV phase (a key parameter necessary to test evolutionary models), the origin and shaping of the LBVNe, and how  the mass-loss behavior (single versus multiple events, bursts) is related to the physical parameters of the central object.  The mass-loss archeology  of the central object can be recovered  from an analysis of its associated nebula. A successful approach is based on a synergistic use of different techniques, at different wavelengths, that allows one to analyze the several emitting components coexisting in the nebula. In particular, a detailed comparison of mid-IR and radio  maps, with comparable spatial resolution, has provided estimates of both ionized gas and dust masses and allowed us to sort out morphological differences in the maps which can be associated with mass-loss behaviour during the LBV phase \\citep{Buemi_2010, Umana_2010}.  \\subsection{The nebula surrounding G79.29+0.46} \\noindent  G79.29+0.46 is  considered a LBV candidate because its observed properties to date do not meet the requirements of spectral and photometric variability to be accepted as a bona-fide member. However, $H_{\\alpha}$ variations, suggestive of mass-loss variability during S-Dor variations, have recently been  reported by \\cite{Vink_08}.  The highly symmetric ringlike structure surrounding G79.29+0.46 was first pointed out by \\citet{Wendker_91}. The thermal nature of the continuum radio emission was determined by \\cite{Higgs_94}, who  suggested that the ringlike nebula is an ionized shell of swept-up interstellar material. By examining the IRAS high-resolution images, \\citet{Waters_96} concluded that the ringlike nebula is a detached shell formed during an epoch of high mass loss ($ \\sim 5 \\times 10^{-4}\\,{\\rm M_{\\odot} yr^{-1}}$) followed by a quieter period. This scenario is consistent with G79.29+0.46 being a LBV, where different mass loss events may have occurred in the recent past.  More recently, the Spitzer Space Telescope has provided high-sensitivity, high-resolution IRAC and MIPS maps \\citep[The Cygnus-X Spitzer legacy program,][]{Hora_2010} which provides a better understanding of the dust properties \\citep{Kraemer_2010}. Moreover, the detection of CO millimetre emission in G79.29+0.46 \\citep{Rizzo_08} demonstrates that another component, consisting of molecular gas, is present in the surroundings of this star and should be included in the budget of the total mass lost from the central object. In their CO maps of G79.29+0.46, \\cite{Rizzo_08} identified components of warm and dense molecular gas whose morphology closely resembles that of the extended IRAS  nebula. They also reported the presence of a shock front, interpreted as a natural consequence of different wind regimes during the central object's evolution.  Despite the wealth of new mid-IR and mm observations of the sources, our knowledge of the radio emission has been limited to the 1988 VLA data reported by \\citet{Higgs_94} and on the 1400 and 350 MHz Westerbork observations, carried out between 1996 and 1997 and reported by \\citet{Setia_03}. In this paper, we present new EVLA observations with sufficient dynamical range, sensitivity, and angular resolution to provide a good match to the Spitzer images. The EVLA observations allow, for the first time, a detailed morphological comparison with other maps tracing the different components of the nebula. ",
        "conclusions": "In this letter we present new radio EVLA observations of the nebula surrounding the LBV G79.29+0.46. In particular,  we have analyzed the 6~cm map and, for the first time, compared the radio free-free, which traces the spatial distribution of the ionized gas, with maps recently obtained in the mid-IR, which traces the spatial distribution of the dust component.  The dust distribution observed at $8\\,\\mu$m appears different than the distribution at $24$ and $70\\,\\mu$m. The warmer dust, traced by the $8\\,\\mu$m emission, is more concentrated in the southwest and southeast regions.  Moreover, when compared with the radio nebula, the $24$ and $70\\,\\mu$m emission appears more extended, while the $8\\,\\mu$m emission is well-contained inside the radio nebula. This is consistent with the existence of two dust components: a cooler component traced by the $24$ and $70\\,\\mu$m emission, and a warmer dust component, probably consisting of smaller grains, traced by the $8\\,\\mu$m emission. This warmer component, with morphological properties quite different from  those observed at longer wavelengths, could be related to the PAH emission that has been claimed by \\citet{Jimenez_2010}.  An improved MIPS $70\\,\\mu$m map, constructed using the most updated pipeline from MIPSGAL team,  detects the presence of a second shell very similar to that seen at $24\\,\\mu$m. Nebular lines and dust continuum are the main possible contributors to the $24\\,\\mu$m emission from objects embedded in a dusty nebula with a hot central component.  The possibility of such a combination, with one kind of emission being predominant, has been suggested by \\citet{Mizuno10} to explain the $24\\,\\mu$m emission detected in the MIPSGAL Bubbles. In the case of G79.29+0.46, the presence of similar shells at $24\\,\\mu$m and at $70\\,\\mu$m, together with the lack of any prominent emission line falling within the response curve of the MIPS $24\\,\\mu$m band,  as evident from the low-resolution IRS spectrum of the shell \\citep{Jimenez_2010}, leaves very little doubt that the shell emission is entirely due to thermal dust emission.  The fact that there are at least two nested shells provides strong constraints on the the origin of the nebula, as it is difficult to explain the presence of multiple shells in the hypothesis of a nebula consisting of swept-up ISM material, as suggested by \\cite{Higgs_94}. It is evident that the dusty shells consist of material ejected by the central object in different mass-loss episodes.  From the $24\\,\\mu$m emission profile, we have determined that the inner shell peaks at $100^{\\prime \\prime}$ and the second one at $200^{\\prime \\prime}$ from the central object. Such dust emission peaks can be related to the epoch when enhanced mass-loss took place. Assuming a distance of 1.7 kpc \\citep{Jimenez_2010}, this corresponds to a linear distance of $0.82$ pc and  $1.64$ pc, respectively. This implies, assuming a shell expansion velocity of $\\sim 30\\,\\rm{km\\,sec}^{-1}$ \\citep{Waters_96}, that the two mass-loss episodes occurred  $2.7 \\times 10^{4}$  and $5.4 \\times 10^{4}$ years ago,  or $1.6  \\times 10^{4}$  and $3.1  \\times 10^{4}$ years ago, if a distance of 1 kpc, as suggested by \\cite{Vink_08}, is used. Our results point out that mass-loss can occur in different episodes, for which we derived characteristics timescales, and pose  quite strong constraints that stellar evolution models must take into account.  The direct comparison of the 6~cm and $24\\,\\mu$m maps (that share comparable spatial resolution) indicates that only part of the inner, brighter nebula is ionized,  (i.e. the nebula is ionization bounded).  The radio nebula shows a sharper profile, depicting regions where probable interaction between the second and the first dusty shells is taking place. This is evident in several regions, most notably in the northeast and in the southwest part of the nebula, where the overall quite regular shell-like morphology is locally disturbed. The mid-IR dust emission depends linearly on density and as a modified blackbody on temperature; the mid-IR emission is therefore most sensitive to dust temperature. On the other hand, assuming an uniform temperature, the thermal radio emission depends on the square of the ionized gas density. This different dependence on density means that the radio emission will emphasize features that have the largest density. This is consistent with the spatial coincidence of the bright radio features with the CO emission reported by \\citet{Rizzo_08} and \\citet{Jimenez_2010}. In particular, the high critical density of the $J=3-2$ CO line ($ n_{3} \\sim 10^{4} \\,\\rm{cm}^{-3}$) indicates that the CO emission traces higher density regions, confirming that the brighter radio features delineate interactions between the shells where shocks can occur."
    },
    "1107/1107.5576_arXiv.txt": {
        "abstract": "The spectra of BL Lac objects and Fanaroff-Riley I radio galaxies are commonly explained by the one-zone leptonic synchrotron self-Compton (SSC) model.  Spectral modeling of correlated multiwavelength data gives the comoving magnetic field strength, the bulk outflow Lorentz factor and the emission region size. Assuming the validity of the SSC model, the Hillas condition shows that only in rare cases can such sources accelerate protons to much above ${10}^{19}$~eV, so $\\gtrsim {10}^{20}$~eV ultra-high-energy cosmic rays (UHECRs) are likely to be heavy ions if powered by this type of radio-loud active galactic nuclei (AGN).  Survival of nuclei is shown to be possible in TeV BL Lacs and misaligned counterparts with weak photohadronic emissions. Another signature of hadronic production is intergalactic UHECR-induced cascade emission, which is an alternative explanation of the TeV spectra of some extreme non-variable blazars such as 1ES 0229+200 or 1ES 1101-232.  We study this kind of cascade signal, taking into account effects of the structured extragalactic magnetic fields in which the sources should be embedded. We demonstrate the importance of cosmic-ray deflections on the $\\gamma$-ray flux, and show that required absolute cosmic-ray luminosities are larger than the average UHECR luminosity inferred from UHECR observations and can even be comparable to the Eddington luminosity of supermassive black holes. Future TeV $\\gamma$-ray observations using the Cherenkov Telescope Array and the High Altitude Water Cherenkov detector array can test for UHECR acceleration by observing $>25$~TeV photons from relatively low-redshift sources such as 1ES 0229+200, and $\\gtrsim$~TeV photons from more distant radio-loud AGN. ",
        "introduction": "Active galactic nuclei (AGN) with extended radio jets powered by super-massive black holes are among the most luminous objects in the low-redshift universe.  Since 2004, when the present generation of imaging atomspheric Cherenkov telescopes began to operate, the number of extragalactic sources detected at $\\gtrsim 0.1 $~TeV (very-high energy; VHE) energies has grown rapidly, and is nearly fifty.\\footnote{See tevcat.uchicago.edu/ and www.mpp.mpg.de/$\\sim$rwagner/sources/}  The \\textit{Fermi} Gamma-ray Space Telescope, now in its fourth mission year, is providing a wealth of new discoveries on $\\gamma$-ray galaxies.  In the high-confidence clean sample of active galactic nuclei associations in the First \\textit{Fermi} LAT AGN catalog \\citep[1LAC;][]{abd101LAC}, more than 600 $\\gamma$-ray blazars, divided about equally into BL Lac objects and flat spectrum radio quasars (FSRQs), were reported.  New classes of GeV $\\gamma$-ray galaxies, e.g., radio-loud narrow line Seyfert galaxies~\\cite{abd09a} and star-forming galaxies powered by supernovae rather than black holes \\citep{abd10sb}, following closely the VHE detections of the starburst galaxies NGC 253~\\citep{ace09} and M82~\\citep{acc09c}, are now firmly established.  Moreover, the spectral energy distributions (SEDs) of radio galaxies detected at GeV and VHE, being misaligned by large ($\\gtrsim 10^\\circ )$ angles to the jet axis and thought to be the parent population of blazars in geometrical unification scenarios \\citep{up95}, are helping to reveal the blazar jet geometry.  About 10 such sources are now detected with \\textit{Fermi}~\\cite{abd10ma}.  Radio-loud AGN detected at TeV energies consist mainly of high-synchrotron-peaked BL Lac objects, including the ultra-variable TeV blazars Mrk 421~\\citep[$z$=0.031;][]{fos08}, Mrk 501~\\citep[$z$=0.033;][]{alb07b} and PKS 2155-304~\\citep[$z$=0.116;][]{aha07a}, and the apparently  non-variable TeV blazars 1ES 0229+200~\\citep{aha07b} and 1ES 1101-232~\\citep{aha07c}. Extragalactic VHE $\\gamma$-ray galaxies include several Fanaroff-Riley (FR) class I radio galaxies~\\citep[Cen A, M87, NGC 1275;][respectively]{aha09a,aha06,ale11c}.  Cen A \\citep{abd10cena}, M87~\\citep{abd09m87}, and NGC 1275~\\citep{abd09ngc1275,2010ApJ...715..554K,BA11} have also been detected at GeV energies.  This list also includes the head-tail radio galaxy IC 310~\\citep{ale10,nsv10}, intermediate-synchrotron-peaked objects like 3C 66A~\\citep{acc09a,ali09,abd11a} and BL Lac~\\citep{alb07a,abd11b}, and the GeV luminous, high-redshift FSRQs 3C 279~\\citep[$z$=0.538;][]{ale11a}, PKS 1510-089~\\citep[$z$=0.361;][]{wag10}, and 4C +21.35~\\citep[PKS 1222+216, $z$=0.432;][]{ale11b}.  The $\\gamma$-ray data from weak-lined BL Lac objects and FR-I radio galaxies are generally well fit with the standard nonthermal electronic synchrotron self-Compton (SSC) relativistic jet model~\\citep[e.g.,][]{mk97,tav98,kat06}, but the use of archival data for highly variable blazars gave large parameter uncertainties in the past.  With the recent simultaneous multi-wavelength data sets for many sources, accurate parameter estimation can be made, either from simple scaling results in the Thomson regime, or from detailed spectral calculations taking into account the Klein-Nishina (KN) effect that is relevant for high, UV/X-ray synchrotron-peaked blazars.  The paper is organized as follows.  In Section 2, we assemble the derived parameter values obtained in various analyses of typical BL Lac objects and FR-I radio galaxies.  From these numbers, we obtain maximum energies of cosmic rays, and show that protons can be accelerated to $\\gtrsim 10$~EeV energies only in a few radio galaxies and flares of BL Lac objects.  Then we explore the associated hadronic signatures expected from TeV blazars in the case where jets of BL Lac objects and FR-I radio galaxies are accelerators of ultra-high-energy cosmic rays (UHECRs, with energies above the ankle of $\\approx {10}^{18.5}$~eV), assuming the validity of the leptonic SSC parameters inferred from rapidly variable AGN. We discuss observable signals produced in the source and calculate those generated outside the source from both $\\gamma$ rays and UHECRs escaping from the jet accelerator and passing through the $\\sim$~Mpc-scale regions of cosmic structure, magnetized clusters and filaments, and the larger $\\sim 100$ Mpc-scale voids of intergalactic space.  For some apparently non-variable TeV blazars, where the one-zone synchrotron/SSC model typically requires extreme parameters for fits, the cascade radiation can be a crucial component of the high-energy radiation spectrum, depending on the strength of the intergalactic magnetic field (IGMF) in voids, as has recently been proposed to explain the $\\gamma$-ray spectra of some extreme blazars such as 1ES 0229+200~\\cite{ek10,ess10,ess11}. We focus on such cascade emissions in the VHE range in Section 3.  Using numerical calculations, we also demonstrate the importance of structured extragalactic magnetic fields (EGMFs) in clusters and filaments for future $\\gamma$-ray detectability by the Cherenkov Telescope Array (CTA) and the High Altitude Water Cherenkov (HAWC) detector array.  Implications of this study and a summary are given in Section 4.  Throughout this work, the cosmological parameters are taken as $H_0=71~{\\rm km}~{\\rm s}^{-1}~{\\rm Mpc}^{-1}$, $\\Omega_m=0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "In this work we studied BL Lac objects and FR-I radio galaxies as potential UHECR sources in light of recent \\textit{Fermi} and imaging atmospheric Cherenkov telescope observations, and considered how future CTA, HAWC, and other high-energy $\\gamma$-ray experiments might test the origin of the $\\gamma$ rays from this class of blazars.  If one accepts the standard synchrotron/SSC model for typical, FR-I galaxies and highly variable BL Lac objects that comprise the majority of VHE radio-loud AGN, the proton maximum energy is typically $\\sim 1-10$~EeV unless UHE protons are produced as rare transient events, and only heavier nuclei normally reach the $\\gtrsim 10^{20}$~eV energies.  In terms of the maximum energy, a heavy-ion dominated composition can be compatible with the standard SSC model because Fe nuclei can be accelerated to $\\gtrsim 10^{20.5}$ eV while surviving against photodisintegration (if $\\delta \\gtrsim 20$; see equation 8). An open issue of the heavy-ion dominated composition scenario of radio-loud AGN is how the significant amount of heavy nuclei is loaded in AGN jets, which is suggested from the PAO composition results~\\cite{abr10} and the observed isotropy in arrival distribution at $\\sim {10}^{19}~{\\rm eV}~Z_{1.5}^{-1} E_{A,20.5}$~\\cite{abr11}.  On the other hand, if hadronic models are adopted for typical, FR-I galaxies and highly variable BL Lac objects, then the observed VHE emission from these objects could be proton synchrotron radiation if protons are accelerated up to $\\sim {10}^{20.5}$~eV, which requires strong magnetic fields, $B' \\sim 10-100$~G that could be found in the inner jets of the radio-loud AGN.  Such hadronic models can be compatible with a proton-dominated composition.  Especially for luminous blazars with spectacular flares and low-peaked BL Lac objects with scattered radiation fields, one may expect high-energy neutrinos produced in inner jets as one of the hadronic signatures~\\cite[e.g.,][]{ad01}.  In either of the synchrotron/SSC or hadronic model, we mainly considered the blazar zone in the inner jet as the emission region of $\\gamma$ rays.  However, some recent studies based on simultaneous radio and $\\gamma$-ray observations are questioning the standard idea that the blazar region is located near the AGN core~\\citep[e.g.,][]{mar10}.  For example, in the case of 3C 345, Schinzel et al. (2012) proposed that the emitting region is located at $\\sim 23$~pc along the jet.  Extreme TeV blazars, in sources like 1ES 0229+200, 1ES 0347-121, H 2346-309, and 1ES 1101-232~\\citep{nv10}, are extreme both in their deabsorbed TeV spectra and their quiescent, non-blazar-like behavior. Their hard source spectra can be explained by $\\gamma$ rays that are produced either via electronic SSC or hadronic processes in inner jets, but hard EBL-deabsorbed VHE spectra typically require extreme source parameters or a special setup~\\cite{tav11,zac11}. An alternative interpretation for the extreme blazars, whose variability is slow or absent, comes from intergalactic cascade emissions.  $\\gamma$-ray-induced or UHECR-induced cascaded emissions can make slowly variable or almost non-variable components provided that the IGMF is weak enough, and the former may be seen as a slowly variable pair-echo component, following a more rapidly variable component.  We examined these possibilities with numerical calculations, taking into account effects of structured EGMFs in filaments and clusters, and demonstrated that the structured EGMFs would play an important role on the proton-induced cascade emission, and that cosmic rays are significantly isotropized especially for the maximum proton energy of $\\lesssim {10}^{19}$~eV. In this case, rather large cosmic-ray luminosities are required in order to explain the emission from extreme TeV blazars such as 1ES 0229+200 by the cascade radiation induced by UHECRs. Note that adopting $E_p^{\\rm max} \\sim {10}^{19}$~eV is motivated by values obtained from the synchrotron/SSC modeling for typical, variable radio-loud AGN rather than extreme, non-variable blazars.  One should keep in mind that the intergalactic cascade scenario itself does not tell about $E_p^{\\rm max}$.  Hence, one may consider protons with $E_p^{\\rm max} \\sim {10}^{20}$~eV, where strong deflections in the structured regions can be avoided.  As a result, the required UHECR luminosity to power VHE emission from extreme TeV blazars becomes more reasonable, which is also comparable to values needed by the proton synchrotron model for typical, variable TeV blazars.  There are some issues in the UHECR-induced cascade scenario that need more discussion. Even if cosmic rays are completely beamed upon entering the voids of intergalactic space, the distant and luminous extreme blazars should be more powerful cosmic-ray sources than the nearby weaker FR-I galaxies found within the GZK radius that are assumed to be responsible for the highest energy UHECRs.  To avoid requiring excessively large beaming-corrected cosmic-ray luminosities to make steady TeV radiation, it is therefore better if the UHECRs from extreme TeV blazars are beamed.  On the other hand, if BL Lac objects and FR-I galaxies are the main sources of UHECRs, then the escaping UHECRs must be significantly isotropized, because most radio galaxies within the GZK radius have misdirected jets.  This isotropization could be caused by radio bubbles or the lobes of radio galaxies.  The contrary behavior may be allowed if the distant extreme blazars from which steady emission is detected are in some ways special, e.g., their host galaxy exists in a region with a very weak EGMF.  Furthermore, note that in the case of rare transient activities with short duration $t_{\\rm dur}$ and longer quiescent periods $t_{\\rm qui} > {\\rm max}[{\\Delta T}_{\\rm CR},{\\Delta t}_{\\rm IGV}]$ between events, the intergalactic cascade scenario for extreme blazars may not work.  The timescale during which cosmic-ray-induced cascade emission lasts for long time is increased by the factor ${\\rm max}[{\\Delta T}_{\\rm CR},{\\Delta t}_{\\rm IGV}]/t_{\\rm dur}$.  Then, the cosmic-ray-induced $\\gamma$-ray flux is consequently reduced by the cosmic-ray and pair-echo induced extension of this emission, and the required cosmic-ray luminosity has to be correspondingly increased, which could make excessive demands on UHECR power.  Such a rare transient episode could, however, produce long-lasting though faint cosmic-ray-induced $\\gamma$-ray emission that lacks associated source emission, in analogous to TeV $\\gamma$-ray-induced cascade radiation~\\citep[e.g.,][]{2010ApJ...719L.130N}. On the other hand, for repeating flaring activities with the time interval $t_{\\rm qui} < {\\rm max}[{\\Delta T}_{\\rm CR},{\\Delta t}_{\\rm IGV}]$, the cosmic-ray-induced cascade emission can be regarded as almost persistent due to contributions from multiple flares.  Observational tests of the properties of the extreme TeV blazars can reveal the radiation mechanism of BL Lac objects and FR-I galaxies and provide a clue to UHECR acceleration, despite the above potential issues.  First, continuing variability searches with VERITAS, HESS, and MAGIC, and future studies with CTA are obviously important to determine if the emission must be made in the jet.  In addition, discrimination of the UHECR-induced intergalactic cascade from the $\\gamma$-ray-induced cascade and attenuated source emission is possible from measurements at $\\gtrsim 1-10$~TeV energies.  Detection of high-energy photons above $25$~TeV from 1ES 0229+200 or above $\\sim$~TeV from more distant blazars, which may be realized by future $\\gamma$-ray detectors such as CTA and HAWC, would be compelling evidence that this kind of object is a source of UHECRs.  In addition to $\\gamma$-ray observations, anisotropy searches with UHECR arrival directions provide another interesting test.  In this work we focused on UHECR acceleration in the inner jets of BL Lac objects and FR-I galaxies.  Note that FSRQs and FR-II galaxies are rarer but more powerful, and they can also be sources of UHECRs and neutrinos. Other scenarios such as shock acceleration at hot spots~\\citep[e.g.,][]{tak90,rb93,th11} and cocoon shocks~\\citep[e.g.,][]{nma95,omy10} are viable for these types of radio-loud AGN.  But their local number density, $n_s \\sim {10}^{-7.5}~{\\rm Mpc}^{-3}$, appears to be too small to avoid strong anisotropy in the local universe~\\cite{ts09}.  In the heavy-ion dominated composition case, not only blazars and radio galaxies, but also radio-quiet AGN~\\cite{pmm09} could be sources of UHE nuclei.  In summary, we have considered observational implications of BL Lac objects and FR-I galaxies as steady sources of UHECRs.  Within the standard synchrotron SSC model for typical, variable BL Lac objects and misaligned counterparts, acceleration of UHE protons to energies $\\gtrsim 10^{20}$~eV is unlikely, so the composition of higher-energy cosmic rays should be dominated by heavy ions within the framework of this model.  The intergalactic cascade emission has to be sub-dominant for highly variable blazars and radio galaxies, while it can play a role on the spectrum of slowly variable or non-variable objects, especially extreme TeV blazars.  If the TeV spectrum of those blazars is produced by UHECR-induced cascade emission, then structured EGMFs, which can significantly isotropize protons, increase the luminosity demands on these sources.  The intergalactic cascade emissions induced by VHE/UHE photons and UHECRs from a distant source can be distinguished by future multi-TeV observations from CTA and HAWC.  In particular, detection of $\\gtrsim 25$~TeV photons from relatively low-redshift sources such as 1ES 0229+200 or $\\gtrsim$~TeV photons from more distant sources would favor such objects as being sources of UHECRs."
    },
    "1107/1107.0165_arXiv.txt": {
        "abstract": "This paper revisits the problem of heat conduction in relativistic fluids, associated with issues concerning both stability and causality. It has long been known that the problem requires information involving second order deviations from thermal equilibrium. Basically, any consistent first-order theory needs to remain cognizant of its higher-order origins. We demonstrate this by carrying out the required first-order reduction of a recent variational model. We provide an analysis of the dynamics of the system, obtaining the conditions that must be satisfied in order to avoid instabilities and acausal signal propagation. The results demonstrate, beyond any reasonable doubt,  that the model has all the features one would expect of a real physical system. In particular, we highlight the presence of a second sound for heat in the appropriate limit. We also make contact with previous work on the problem by showing how the various constraints on our system agree with previously established results. ",
        "introduction": "Relativistic thermodynamics continues to provide interesting challenges, in particular in the context of dissipative and nonlinear phenomena. The issues involved range from direct applications in various areas of physics to fundamental problems like the nature of time (visavi the second law of thermodynamics) and the formation of structures at nonlinear deviations from thermal equilibrium. Much recent work has been motivated by the modelling of complex astrophysical systems, like neutron stars \\cite{livrev}, and cosmology \\cite{maartens}. There has also been a resurgence of interest in dissipative systems in the context of colliders like RHIC at Brookhaven and the LHC at CERN \\cite{rhic1,rhic2,rhic3}. These latter developments, which have to a large extent been driven by the need to understand the dynamics of a hot quark-gluon plasma, are often linked with underlying principles like the AdS/CFT conjecture and holography \\cite{rangamani}. Even though the problem dates back to the origins of relativity theory, it remains (in a slightly different guise) at the forefront of modern thinking.  According to the established consensus view, one must account for second-order deviations from thermal equilibrium in order to achieve causality and stability.  This is certainly the lesson from the celebrated work of Israel and Stewart \\cite{IS,IS2}, see \\cite{isrec1,isrec2,isrec3} for recent work on the problem. We have recently revisited the key points in the context of heat conduction \\cite{cesar}, taking a multi-fluid prescription based on  Carter's convective variational formulation for relativistic fluids \\cite{carter02} as our starting point. This is a mathematically elegant approach that has the flexibility required to account for the physics that we need to consider. A particularly appealing feature of the variational approach is that, once an ``equation of state'' for matter is provided, the theory provides the relation between the various currents and their conjugate momenta \\cite{livrev}. The variational analysis leads to a second-order model which has the key elements required for causality and stability, in particular, it clarifies the role of the inertia of heat (e.g., the effective mass associated with phonons). This effect enters the model in an intuitive fashion in terms of entrainment between the matter and heat \\cite{nilsclass}. As demonstrated by Priou \\cite{priou} some time ago, the final variational model is formally equivalent to the Israel-Stewart construction. This exercise demonstrates clearly that the relaxation associated with causal heat transport is determined by the thermal inertia. At the end of the day, the theoretical framework becomes rather intuitive and the physics involved seems natural.  Does this mean that no troublesome issues remain in this problem area? Not quite. First of all, it is clear that the need to introduce additional parameters (e.g., the relevant relaxation times) and keep track of higher order terms (fluxes of the fluxes etcetera) make actual applications rather complex. Secondly, we are not much closer to considering systems that deviate significantly from equilibrium, such that there is no natural ``small'' parameter to expand in. The variational model sheds some light on this regime by clarifying the role of the temperature in systems out of equilibrium, but there is some way to go before we understand issues associated with, for example, any ``principle of extremal entropy production'' and instabilities that lead to structure formation. Finally, despite the obvious successes of the extended thermodynamics framework \\cite{joubook}, there is no universal agreement concerning the validity (and usefulness) of the results. To some extent this is natural given the interdisciplinary nature of the problem; to make progress we need to account for both thermodynamical principles and fundamental general relativity. This leads to a range of deep questions concerning, in particular, the actual meaning of the variables involved in the different models (e.g., the entropy). The ultimate theory must have a clear link with statistical physics and even information theory. Our efforts are not yet at that level. Basically, we need to continue to make progress if we are to address fundamental problems in, for example, cosmology.  This paper sets a rather more modest target; we want to explore the extent to which a ``first-order'' formulation for heat conduction in general relativity is viable. The question may seem somewhat odd given that we have already acknowledged the need to account for (at least) second order contributions. However, it is interesting to ask whether a first-order model may nevertheless be useful (possibly in a somewhat restricted sense). We will demonstrate that this is, indeed, the case. Noting that the original first order models, due to Eckart \\cite{eckart} and Landau and Lifshitz \\cite{landau}, were incomplete we develop a consistent framework that includes the key thermal relaxation. We then consider the properties of this model, and show that it can be made both stable and causal (making contact with the classic work by Hiscock and Lindblom \\cite{hl1,hl2}). This does not mean that the system may not exhibit instabilities. On the contrary, instabilities are in a sense generic in these problems \\cite{2stream}, but the analysis sheds further light on the nature of these instabilities and also elucidates the stabilizing role of the thermal inertia. The discussion also provides insight into the emergence of second sound (an effect that has been experimentally verified in low temperature crystals) associated with the heat transport. This provides a key link to systems that exhibit superfluidity, and demonstrates the potential for a unified treatment of heat transport in normal and superfluid matter. ",
        "conclusions": ""
    },
    "1107/1107.4093_arXiv.txt": {
        "abstract": "{We present the results of a study of weak gravitational lensing by galaxies using imaging data that were obtained as part of the second Red Sequence Cluster Survey (RCS2). In order to compare to the baryonic properties of the lenses we focus here on the $\\sim$300 square degrees that overlap with the data release 7 (DR7) of the Sloan Digital Sky Survey (SDSS). The depth and image quality of the RCS2 enables us to significantly improve upon earlier work for luminous galaxies at $z\\geq0.3$. To model the lensing signal we employ a halo model which accounts for the clustering of the lenses and distinguishes between satellite and central galaxies. Comparison with dynamical masses from the SDSS shows a good correlation with the lensing mass for early-type galaxies. The correlation is less clear for late-type galaxies, possibly due to rotation. For low luminosity (stellar mass) early-type galaxies we find a satellite fraction of $\\sim$40\\% which rapidly decreases to $<10$\\% with increasing luminosity (stellar mass). The satellite fraction of the late-types has a value in the range 0-15\\%, independent of luminosity or stellar mass. At high masses the satellite fraction is not well constrained, which we partly attribute to the modelling assumptions. To infer virial masses we apply simple models based on an independent satellite kinematics analysis to account for intrinsic scatter in the scaling relations. We find that early-types in the range $10^{10}<L_r<10^{11.5} L_{\\odot}$ have virial masses that are about five times higher than those of late-type galaxies and that the mass scales as $M_{200} \\propto L^{2.34^{+0.09}_{-0.16}}$. For an early-type galaxy with a fiducal luminosity of $10^{11}L_{r,\\odot}$, we obtain a mass $M_{200}=(1.93^{+0.13}_{-0.14})\\times10^{13}h^{-1}M_{\\odot}$. We also measure the virial mass-to-light ratio, and find for $L_{200}<10^{11}L_{\\odot}$ a value of $M_{200}/L_{200}=42\\pm10$ for early-types, which increases for higher luminosities to values that are consistent with those observed for groups and clusters of galaxies. For late-type galaxies we find a lower value of $M_{200}/L_{200}= 17\\pm9$. Our measurements also show that early- and late-type galaxies have comparable halo masses for stellar masses $M_*<10^{11}M_{\\odot}$, whereas the virial masses of early-type galaxies are higher for higher stellar masses. To compare the efficiency with which baryons have been converted into stars, we determine the total stellar mass within $r_{200}$. Our results for early-type galaxies suggest a variation in efficiency with a minimum of $\\sim$10\\% for a stellar mass $M_{*,200}=10^{12}M_{\\odot}$. The results for the late-type galaxies are not well constrained, but do suggest a larger value.} ",
        "introduction": "\\hspace{4mm} There is now overwhelming evidence that galaxies are surrounded by dark matter haloes. Studying the global properties of the haloes, such as their virial masses or density profiles, however, has proven difficult due to a lack of reliable tracers of the gravitational potential at large distances. Improving observational constraints is important because the details of galaxy formation are not completely clear, even though significant progress has been made in recent years \\citep[e.g.][]{Bower10,Kim09}. The relation between the baryons and the dark matter in galaxies has been studied using numerical simulations \\citep[e.g.][]{Wang06,Croton06,Somerville08,Moster10} and it is important to confront the predictions with observations. This requires reliable estimates of both the dark matter and the baryonic content of galaxies. \\\\ \\indent Several observables can be used to trace the baryons, such as the luminosity of a galaxy, which is readily available. It is also possible to derive stellar masses by fitting stellar synthesis models to either the spectral features of a galaxy \\citep{Kauffmann03,Gallazzi05} or to its colours \\citep{BelldJ01,Salim07}. The stellar mass estimates are tightly correlated to various other important global properties of galaxies \\citep[colour, metallicity, luminosity, environment, see e.g.][and references therein]{Grutzbauch11} and they are therefore considered a useful tracer of the baryonic content of a galaxy. \\\\ \\indent Numerical simulations suggest that the dark matter haloes of massive galaxies extend out to hundreds of kiloparsecs \\citep[e.g.][]{Springel05}, which is supported by observations \\citep[e.g.][]{Hoekstra04}. For nearby galaxies it is possible to study the dark matter distribution using the dynamics of planetary nebulae \\citep[e.g.][]{Napolitano09}. In addition, studies of satellite galaxies around central galaxies \\citep[e.g.][]{More11,Conroy07} have provided constraints on the relation between baryons and dark matter. Unfortunately these studies require spectroscopy of large numbers of objects, which makes them rather expensive. Furthermore, the observations are limited to small scales due to the requirement of having optical tracers, which complicates the determination of the virial mass of the haloes galaxies reside in, unless one is willing to extrapolate the measurements. \\\\ \\indent Fortunately it is possible to probe the matter distribution on large scales, thanks to an effect called weak gravitational lensing; we can measure the distortion of the shapes of faint background galaxies (sources) caused by the bending of light rays by intervening mass concentrations (lenses). The distortion is independent of the type of matter in the lenses, and so the projected mass of the lens is measured without any assumption on the physical state of the matter at scales from a few kiloparsec to a few megaparsec.  \\\\ \\indent The weak lensing signal around a single galaxy is too weak to detect since it is 10-100 times smaller than the intrinsic ellipticities of galaxies. Therefore the galaxy-galaxy signal has to be averaged over many lenses to decrease the shape noise. Although individual galaxies cannot be studied in this way, their average properties can be determined \\citep[e.g.][]{Brainerd96,Fischer00,Hoekstra04}. Only more recently has it become possible to study lenses as a function of properties such as type, luminosity, stellar mass, etc., because early studies lacked the ancillary data needed to subdivide the lenses into subsamples. For instance \\citet{Hoekstra05} used nearly 34 square degrees of the Red Sequence Cluster Survey (RCS) \\citep{Gladders05} for which photometric redshifts were available \\citep{Hsieh05}, to study the relation between the virial mass and baryonic contents of isolated galaxies in the redshift range $0.2<z<0.4$, and derived star formation efficiencies for early- and late-type galaxies. Thanks to the wealth of ancillary data, the Sloan Digital Sky Survey \\citep[SDSS;][]{York00} has had a major impact on galaxy-galaxy lensing studies \\citep[e.g.][]{GuzikS02,Mandelbaum06}. This is evidenced by \\citet{Mandelbaum06} who used nearly 5000 square degrees of the SDSS DR4 \\citep{Adelman06} to study galaxies in the redshift range $0.02<z<0.35$ as a function of galaxy type and environment, and constrained the stellar mass to virial mass relation, the luminosity to virial mass relation and the satellite fractions of the lens samples. \\\\ \\indent Currently no survey can surpass the precision that can be achieved by the SDSS at low redshift ($z<0.3$) because of the large survey area and the availability of spectroscopic data. We note, however, that complementing the SDSS data with deeper imaging by the Panoramic Survey Telescope \\& Rapid Response System\\footnotemark \\footnotetext[1]{http://pan-starrs.ifa.hawaii.edu/public/} \\citep[Pan-STARRS;][]{Kaiser02} will provide a major improvement, as is demonstrated by the results we present here. For lenses with $z>0.3$ it is possible to achieve a significant improvement over the SDSS results by surveying a smaller area with deeper data and good image quality; it allows us to use sources at higher redshifts. This is important because the amplitude of the lensing signal scales proportionally to the ratio of the angular diameter distance between the lens and the source and the distance between the observer and the source. The signal decreases rapidly when the lens redshift approaches the peak of the source redshift distribution, which occurs around $z\\sim 0.35$ for the SDSS. \\\\ \\indent In this paper we use data from the second generation Red Sequence Cluster Survey (RCS2) to measure the weak lensing signal around galaxies that are observed in the SDSS. The RCS2 is a nearly 900 square degree imaging survey carried out by the Canada-France-Hawaii-Telescope (CFHT), and is $\\sim$2 magnitudes deeper than the SDSS in $r'$. The increase in depth combined with a median seeing of 0.7\", which is a factor of two smaller than the seeing in the SDSS, results in a source galaxy number density that is about five times higher, and a source redshift distribution that peaks at z$\\sim$0.7.  \\\\ \\indent We use the overlapping area between the two surveys, which amounts to approximately 300 square degrees, in order to assign the spectroscopic redshifts, luminosities, stellar masses and dynamical masses from the SDSS to the lenses. The lensing analysis itself is performed on the RCS2 data. Even though the overlap between the surveys is modest, the loss in survey area is outweighed by the gain in the number density of source galaxies and the improvement of the lensing efficiency. This enables us to improve the measurements of the lensing signal around the most massive galaxies, which mostly reside at redshifts where the SDSS is not very sensitive. \\\\ \\indent In this paper we describe the lenses in Section \\ref{sec_lenssamp}. The weak lensing analysis is discussed in Section \\ref{sec_lensing}. The halo model that we have implemented is introduced in Section \\ref{sec_halo}. In Section \\ref{sec_mdyn} we compare the weak lensing mass to the dynamical mass. We describe the luminosity results in Section \\ref{sec_lumi}, and the stellar mass results in Section \\ref{sec_mstel}. We summarize our conclusions in Section \\ref{sec_conc}. Throughout the paper we assume a WMAP5 cosmology \\citep{Komatsu09} with $\\sigma_8=0.8$, $\\Omega_{\\Lambda}=0.73$, $\\Omega_M=0.27$, $\\Omega_b=0.045$ and the dimensionless Hubble parameter $h=0.7$. All distances quoted are in physical (rather than comoving) units unless explicitly stated otherwise. ",
        "conclusions": "} We measured the halo masses for early- and late-type galaxies and compared these to their luminosity and stellar mass. For this purpose, we measured the weak lensing signal induced by the galaxies with SDSS spectroscopy that overlap with the RCS2, and modelled the data with a halo model. This enabled us to improve the constraints on the lensing measurements for the most massive galaxies, which typically reside at redshifts where the SDSS is not very sensitive.  \\\\ \\indent The halo mass and the dynamical mass correlate well for early-type galaxies, but not for late-type galaxies. A likely explanation is that late-type galaxies are rotating, resulting in an overestimation of the velocity dispersion, and hence of the dynamical mass. Furthermore, in contrast to the dynamical mass, the weak lensing mass can easily be related to numerical simulations, and provides constraints for the models that describe the relationship between baryons and dark matter. \\\\ \\indent The halo masses of galaxies increase with luminosity and stellar mass. For a given luminosity, the halo mass of the early-types is on average about five times larger than the late-types. We fitted a powerlaw relation between the luminosity and halo mass, and find that in the range $10^{10}<L_r<10^{11.5} L_{\\odot}$, the halo mass scales with luminosity as $M_h \\propto L^{2.34^{+0.09}_{-0.16}}$ and $M_h \\propto L^{2.2^{+0.7}_{-0.6}}$ for the early- and late-type galaxies respectively. For an early-type galaxy with a fiducal luminosity $L_0=10^{11}L_{r,\\odot}$, we obtain a mass $M_{200}=(1.93^{+0.13}_{-0.14})\\times10^{13}h^{-1}M_{\\odot}$. We computed $L_{200}$, the additional luminosity around the lenses within $r_{200}$, and find that the $M_{200}/L_{200}$ ratio of the early-types is larger than for the late-types: for $L_{200}<10^{11}L_{\\odot}$ we find $M_{200}/L_{200}=42\\pm10$ for early-types, whilst $M_{200}/L_{200}= 17\\pm9$ for late-types. This suggests that the difference in halo mass is not solely due to the fact that early-types reside in denser environments, but is at least partly intrinsic. \\\\ \\indent Below a stellar mass of $10^{11}M_{\\odot}$ the halo mass of early- and late-types are comparable. For larger stellar masses, the best fit halo masses of the early-types are larger than the late-types. We computed $M_{*,200}$, the total stellar mass within $r_{200}$, in order to calculate the baryon conversion efficiency $\\eta$. Our results for early-type galaxies suggest a variation in efficiency with a minimum of $\\sim$10\\% for a stellar mass $M_{*,200}=10^{12}M_{\\odot}$. The results for the late-type galaxies are not well constrained, but do suggest a larger value. \\\\ \\indent The satellite fraction is $\\sim$40\\% for the low luminosity (stellar mass) early-type galaxies, and decreases rapidly to $<10$\\% with increasing luminosity (stellar mass). The satellite fraction of the late-types has a value in the range 0-15\\%, independent of luminosity or stellar mass. The satellite fraction is difficult to constrain at the high stellar mass/luminosity end, as the shape of the combined shear signal from the satellites mimics an NFW profile. Decreasing the truncation parameter leads to tighter constraints, and appears to be justified for the most massive early-type satellites based on the N-body simulation results of \\citet{Limousin09}. Additional support comes from studying the shear signal of massive early-type galaxies that were selected to be satellites shown in Appendix \\ref{ap_sat}, but the errors are currently too large to constrain the fraction of dark matter that is stripped. A more realistic description of the stripping of the haloes of massive satellite galaxies may result in an improvement of the constraints on the satellite fraction from weak lensing studies alone. \\\\ \\indent The halo mass appears to decrease with redshift for the highest stellar mass bins, a trend that is qualitatively in agreement with predictions from numerical simulations. The signal-to-noise on the measurements is currently too low to provide a detailed view on the growth of dark matter haloes, but it shows that with future surveys weak lensing can be used to study in great detail the evolution of the relation between baryons and dark matter.  \\paragraph{Acknowledgements \\\\ \\\\} We would like to thank Jarle Brinchman for his extensive help with the MPA/JHU catalogues, Alexie Leauthaud for useful discussions and suggestions concerning the halo model, and Thomas Erben and Hendrik Hildebrandt for their help with the THELI pipeline. Additionally, we would like to thank our colleagues at Leiden Observatory, in particular Konrad Kuijken and Elisabetta Semboloni, and Tim Schrabback, for useful suggestions and comments. Finally, we would like the anonymous referee for useful comments. \\\\ \\indent H.H. acknowledges support from a Marie Curie International Reintegration Grant, and from a VIDI grant from the Nederlandse Organisatie voor Wetenschappelijk Onderzoek (NWO). The RCS2 project is supported in part by grants to H.K.C.Y. from the Canada Research Chairs program and the Natural Science and Engineering Research Council of Canada. \\\\ \\indent This work is based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institute National des Sciences de l'Univers of the Centre National de la Recherche Scientifique of France, and the University of Hawaii. We used the facilities of the Canadian Astronomy Data Centre operated by the NRC with the support of the Canadian Space Agency."
    },
    "1107/1107.1057.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We aim to constrain the location, composition, and dynamical state of planetesimal populations and dust around the young, sun-like (G2\\,V) star \\object{HD\\,107146}.} % methods heading (mandatory) {We consider coronagraphic observations obtained with the Advanced Camera for Surveys ({\\it HST}/ACS) onboard the {\\it HST} \\normalfont in broad V ($\\lambda_{c}\\approx 0.6\\,\\mu\\rm{m}$) and broad I ($\\lambda_{c}\\approx 0.8\\,\\mu\\rm{m}$) filters, a resolved 1.3\\,mm map obtained with the {\\it Combined Array for Research in Millimeter-Wave Astronomy} \\normalfont ({\\it CARMA})\\normalfont, {\\it Spitzer}/IRS low resolution spectra in the range of $7.6\\,\\mu\\rm{m}$ to $37.0\\,\\mu\\rm{m}$, and the spectral energy distribution (SED) of the object at wavelengths ranging from $3.5\\,\\mu\\rm{m}$ to $3.1\\,\\rm{mm}$. We complement these data with new coronagraphic high resolution observations of the debris disk using the Near Infrared Camera and Multi-Object Spectrometer (\\itshape{HST}\\normalfont/NICMOS) aboard the \\itshape{Hubble Space Telescope} \\normalfont (\\itshape{HST}\\normalfont) in the F110W filter ($\\lambda_{c}\\approx 1.1\\,\\mu\\rm{m}$). The SED and images of the disk in scattered light as well as in thermal reemission are combined in our modeling using a parameterized model for the disk density distribution and optical properties of the dust.} % results heading (mandatory) {A detailed analytical model of the debris disk around HD\\,107146 is presented that allows us to reproduce the almost entire set of spatially resolved and unresolved multi-wavelength observations. Considering the variety of complementary observational data, we are able to break the degeneracies produced by modeling SED data alone. We find the disk to be an extended ring with a peak surface density at 131\\,AU. Furthermore, we find evidence for an additional, inner disk probably composed of small grains released at the inner edge of the outer disk and moving inwards due to Poynting-Robertson drag. A birth ring scenario (i.e., a more or less broad ring of planetesimals creating the dust disk trough collisions) is found to be the most likely explanation of the ringlike shape of the disk.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "\\label{intro}  Debris disks were first discovered via infrared excess emission associated with main sequence stars such as Vega from observations with the {\\it Infrared Astronomical Satellite} \\citep[{\\it IRAS};][]{aum84}. The first spatially resolved image of such an object was obtained around $\\beta$\\,Pictoris \\citep{smi84}. In the last few years the {\\it Spitzer Space Telescope} has increased the number of known debris disk candidates identified by their thermal emission above photospheric levels to several hundreds. As the most readily detectable signposts of other planetary systems, debris disks help us to improve our understanding of the formation and evolution of them as well as of our own solar system \\citep{mey07, wya08, kri10}. The dust detected in such disks is thought to be removed from those systems by the stellar radiation on short time scales compared to their ages. Thus, it must be transient, or more likely continuously replenished by ongoing collisions of bigger objects like planetesimals left over from the planet formation process. For isolated systems the circumstellar dust's spatial distribution can be influenced by gravitational interaction with planets in addition to the influence of the star \\citep{dom03, ken04, str06, wya08}. However, studying the spectral energy distribution (SED) of the dust alone provides only weak, ambiguous constraints for its properties such as chemical composition, grain size, and spatial distribution \\citep{wol03}. For example, the location of the inner disk radius of the dust distribution is strongly degenerate with the lower limit of the grain size distribution. Direct measurements of the spatially resolved surface brightness by imaging mitigate many of these degeneracies. Images of disks seen near to face-on are particularly useful to directly determine the azimuthal and radial distribution of the dust and thus to trace the density structure.  \\begin{figure*} \\centering \\includegraphics[width=1\\linewidth]{imaSca.pdf} \\caption{Scattered light images. The {\\it HST}/ACS images at $0.6\\,\\mu\\rm{m}$ (left) and at $0.8\\,\\mu\\rm{m}$ (center) and the {\\it HST}/NICMOS image at $1.1\\,\\mu\\rm{m}$ (right). The images are oriented north-up and the {\\it HST}/ACS images are rebinned to the {\\it HST}/NICMOS scale of $0.076\\,\\rm{\\arcsec/pix}$. The plotted coordinate system indicates the distance from the star in arcseconds along the two axes. Note the different flux scales in $\\rm{mJy/as}^{2}$ and the fact that the stellar PSF residuals are heavily truncated in these scales. The object in the SW is a background galaxy.} \\label{imaSca}% \\end{figure*}  In this paper, we present new spatially resolved observations of the HD\\,107146 debris disk obtained with the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) onboard the {\\it Hubble Space Telescope} ({\\it HST}). These data complement existing spatially resolved images at optical and millimeter wavelengths as well as spatially unresolved data of the star and the disk over a broad range of wavelengths (Table~\\ref{SED_data}). An approach to combine all available data and so to derive a detailed model of the disk and to break degeneracies produced by modeling single datasets is presented.  HD\\,107146 is a young, well studied, solar-type star (G2\\,V), $28.5\\pm0.7\\,\\rm{pc}$ \\citep{per97} from the sun. Its luminosity and effective temperature are $\\approx 1.1\\,\\rm{L_{\\sun}}$ \\citep{ard04} and $\\approx 5850\\,\\rm{K}$, respectively. Its estimated age ranges from 80 to 320\\,Myr \\citep{wil04, moo06, roc09}. \\citet{sil00} was the first to note the presence of an IR excess based on {\\it IRAS} measurements. The origin of this excess in a debris disk system was discussed by \\citet{met04}. HD\\,107146 has been well studied in the context of the {\\it Spitzer} Legacy Science Program {\\it FEPS} \\citep[Formation and Evolution of Planetary Systems;][]{mey06} resulting in a broad variety of photometric measurements as well as spectra of the star and the disk \\citep[e.g.][]{hil08}. The debris disk around HD\\,107146 is remarkable in many respects. With its large inner hole of $\\sim 60\\,\\rm{AU}$ and its extension of $> 200\\,\\rm{AU}$ in radius, corresponding to $2\\arcsec$ and $7\\arcsec$, respectively \\citep{ard04}, it can be imaged by coronagraphy without masking large parts of the disk. Its particularly high scattered light surface brightness enables high signal-to-noise imaging \\citep{ard04}. The disk is seen with an inclination of $25^\\circ \\pm 5^\\circ$ from face-on \\citep{ard04}. This means, that the radial surface brightness distribution can be observed without strong projection effects which would occur in a close to edge-on case. With the inclination known, one can properly account for the remaining effects. In addition, the inclination of the disk enables us to constrain the dust properties from potential scattering phase asymmetries.  We present our new observations in Sect.~\\ref{currObs}. The available data collected from the literature is summarised in Sect.~\\ref{avData}. Details about our modeling, particularly the modeling software, the applied model, the order in which we include the available data and the information we extract in each step, respectively, can be found in Sect.~\\ref{modeling}. The results are presented and discussed in detail in Sect.~\\ref{discussion}. Conclusions are given in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  A detailed model of the debris disk around HD\\,107146 has been created and fitted to all available data using the approach described in the paper. A broad range of resolved data from optical scattered light to millimeter data as well as the SED has been reproduced by a single radial density profile. All parameters of our model have been fixed without remaining degeneracies using only one scattered light image and the well-sampled SED of the disk in the particular case (Sect.~\\ref{modeling}). We found the disk to be a broad ring of dust with smooth inner and outer boundaries, and with a peak position from the star and FWHM of the surface density of $\\approx 131\\,\\rm{AU}$ and $\\approx 91\\,\\rm{AU}$, respectively. We concluded from our modeling results that the disk is heavily collision dominated. Discrepancies of our model to the data are discussed. From these discrepancies we find strong evidence for an additional inner disk component, possibly near the habitable zone (the stellocentric distance where liquid water can exist on planetary surfaces, i.e., few AU from a solar-type star) of the star and for a different chemical composition or shape of the dust grains. The additional inner disk component is predominantly composed of small grains and the results from our modeling are consistent with these grains being released at the inner edge of the disk and dragged inwards due to Poynting-Robertson drag. Only detailed, simultaneous modeling of the large amount of available high quality, complementary data was able to reveal these signposts. We did not find evidence of an orbiting planet from the available data and found that a birth ring scenario is likely responsible for the ringlike shape of the disk."
    },
    "1107/1107.3399_arXiv.txt": {
        "abstract": "{Using Milky Way data of the new Effelsberg--Bonn \\ion{H}{i} Survey (EBHIS) and the Galactic All-Sky Survey (GASS), we present a revised picture of the high-velocity cloud (HVC) complex Galactic center negative (GCN). Owing to the higher angular resolution of these surveys compared to previous studies (e.g., the Leiden Dwingeloo Survey), we resolve complex GCN into lots of individual tiny clumps, that mostly have relatively broad line widths of more than $15\\,\\mathrm{km\\,s}^{-1}$. We do not detect a diffuse extended counterpart, which is unusual for an HVC complex. In total 243 clumps were identified and parameterized which allows us to statistically analyze the data.  Cold-line components (i.e., $\\Delta v_\\mathrm{fwhm}<7.5\\,\\mathrm{km\\,s}^{-1}$) are found in about 5\\% only of the identified cloudlets. Our analysis reveals that complex GCN is likely built up of several subpopulations that do not share a common origin. Furthermore, complex GCN might be a prime example for warm-gas accretion onto the Milky Way, where neutral \\ion{H}{i} clouds are not stable against interaction with the Milky Way gas halo and become ionized prior to accretion.} ",
        "introduction": "In recent years it has become more and more clear that the gas in the Milky Way (MW) halo plays an important role for understanding the evolution of our host galaxy. Matter is constantly expelled from the Galactic disk via winds and supernovae-driven outflows and at the same time material is accreted to the disk. Furthermore, gas produced by tidal interaction or remnants from the merging history of the MW may still reside in the halo \\citep[e.g.,][ and references therein]{sancisi08}. Today, it is known that the so-called intermediate- and high-velocity clouds and complexes represent various stages of these processes \\citep[e.g.,][ and references therein]{kalberla09}. A prime example for tidally disrupted material is the Magellanic Stream (MS) that extends over hundreds of degrees over the sky. In many other cases the origin of the gas is still debated.  One of these cases is the high-velocity cloud (HVC) complex Galactic center negative (GCN) which is one of the smaller cloud complexes \\citep[in total flux and spatial size;][]{wakker91}. Owing to the lack of suitable background sources (stars or QSOs), relatively little is known about its chemical composition, gas properties, and distance.  There are several early studies of complex GCN. The first \\ion{H}{i} detection of an individual cloud belonging to GCN was described in \\citet{saraber74}. This object is located close to the Galactic center ($(l,b)=(8\\degr,-4\\degr)$). Several origin scenarios were discussed in Saraber \\& Shane's work, (1) an extra-galactic origin (with the cloud being a galaxy), (2) that it is an object penetrating the MW disk, and (3) that it is ejected from the Galactic center. In any case, the cloud would be located at a distance of several kpc (using some simplifications). The same cloud was independently studied by \\citet{mirabel76}. They arrive at the conclusion that the object was ejected from the Galactic center and exclude other potential scenarios, e.g., that it is an extension of the Magellanic Stream, which is clearly contradicting \\citet{giovanelli81}, who favored the MS scenario. Another GCN cloud was observed by \\citet{cohen78}. Again, several origins were discussed, with an additional scenario coming into play, i.e., that the cloud is a local group object (but not a galaxy). \\begin{figure*} \\centering \\includegraphics[width=0.98\\textwidth,bb=14 40 691 532,clip=]{fig1.eps} \\caption{Column density map of HVC complex GCN. Spurious emission is visible along the MW plane that is caused by the increased noise level in this direction. The column densities were computed over the velocity interval $v_\\mathrm{lsr}=-360\\ldots-160\\,\\mathrm{km\\,s}^{-1}$ using a mask to suppress image artifacts caused by noise. The mask was created using a hanning filtered data cube (hanning-9) and applying a $2.5\\sigma$ threshold. The graticule shows Galactic coordinates.} \\label{figclumpcolumndensitymap}% \\end{figure*}  Even the most recent analyses of the neutral hydrogen in GCN are based on observations dating back to about three decades ago \\citep{bajaja85,hulsbosch88} which were performed with the Instituto Argentino de Radioastronomia 30-m and the Dwingeloo radio telescope. Both surveys were not fully sampled and hence are missing information. Furthermore, the small dish sizes led to a poor angular resolution of about $30\\arcmin$. In these data sets GCN appears to be a population of widely scattered, small clouds at high negative velocities with sizes between 1 and 15 deg$^2$ \\citep{wakker91}.  Most authors place GCN at a distance of more than about 10\\,kpc such that the size of the complex is at least on the same order \\citep[e.g.,][]{saraber74,wakker97,weiner01,wakker04} though the uncertainties are high because no direct distance estimates are available.  Using the X-ray shadowing technique, \\citet{kerp99} could show that in the line of sight toward the galaxy Mrk\\,509 (which is located in the GCN footprint) there is excess soft X-ray emission. For the same sight line \\citet{sembach95} performed HST absorption-line measurements and detected highly ionized high-velocity gas. The X-ray data suggest that there is collisionally ionized gas, while \\citet{sembach95} may have detected the cooler portion of the collisionally ionized plasma.   A few additional sight lines revealed highly ionized gas  \\citep[e.g., \\ion{C}{iv} and \\ion{Si}{iv},][and references therein]{wakker01}. The presence of these metals provides strong evidence that GCN is not of primordial origin. Some authors conclude that GCN represents an inflow of gas to the Galactic center \\citep{mirabel82,mirabel84} as well as the complex Galactic center positive (GCP), containing the famous Smith Cloud \\citep[see][for a recent study]{lockman08}. Based on Leiden/Argentine/Bonn survey data \\citep[LAB;][]{kalberla05}, \\citet{jin10} used several clouds in the direction of complex GCN and calculated a hypothetical orbit that these objects might follow concluding that GCN likely was tidally stripped of a dwarf galaxy.   In this work we present new results obtained using the recently completed Galactic All-Sky Survey \\citep[GASS;][]{mcclure2009,kalberla2010} obtained with the Parkes telescope and data from the ongoing Effelsberg--Bonn \\ion{H}{i} Survey \\citep[EBHIS;][]{winkel10a,kerp11} measured with the 100-m telescope at Effelsberg. The better angular resolution compared to previous large-area \\ion{H}{i} surveys  and the full sampling of the data allowed us to obtain a revised picture of complex GCN.   The paper is organized as follows. In Section\\,\\ref{secdata} we describe the data from EBHIS and GASS that were used for the analysis. We compiled a catalog of clouds (Section\\,\\ref{seccatalog}) to perform a statistical analysis, the results of which are presented in Section\\,\\ref{secstatistics}. In Section\\,\\ref{secdiscussion} we discuss the results and conclude with a summary and outlook (Section\\,\\ref{secsummary}). ",
        "conclusions": "\\label{secdiscussion} The observed properties of clouds in our sample show that GCN is quite different compared to other large HVC complexes. Not only the small spatial size of most of the clumps and the low fraction of cold-line components are rather untypical, but also the missing diffuse extended \\ion{H}{i} emission, which is detected in most other HVC complexes \\citep[e.g.,][using the less sensitive LAB survey data]{kalberla06}. This is especially true because the absorption line data indicate the presence of a (low-column density) diffuse gas phase.  Furthermore, we find many signs for collisionally or photo-ionized gas in the region. Several head--tail features were observed, and we find a decrease of the maximal \\ion{H}{i} column densities toward lower absolute radial velocities. This makes it likely that GCN clouds are subject to interaction processes with an ambient medium (i.e., the Milky Way halo).  Unfortunately, the large distance uncertainty does not allow us to estimate important physical parameters like temperature or pressure with sufficient accuracy to rule out or accept the various possible interaction processes directly, i.e. Rayleigh--Taylor- and Kelvin--Helmholtz instabilities, ram-pressure, or shocks (which could also lead to fragmentation). However, we will discuss below some of the mechanisms under the assumption that clouds in the N--S stripe are indeed related to the Leading Arm, as suggested by the simulations of \\citet{diaz11}, which constrains the distance to 14 to 22\\,kpc. This is at least consistent with earlier (rough) distance estimates and partially with the orbit proposed by \\citet{jin10}.  In this case, for a beam-sized (unresolved) cloudlet at 18\\,kpc distance with a typical column density of $10^{19}\\,\\mathrm{cm}^{-2}$ we estimate the volume density to be on the order of $n_\\mathrm{c}=0.1\\,\\mathrm{cm}^{-3}$. The size of such an object would be about 50\\,pc. Furthermore, using the MW model of \\citet{kalberla07}, we can obtain an estimate of the pressure and density of the surrounding halo gas along the N--S filament: $P/k_\\mathrm{B}\\sim10^3\\,\\mathrm{K\\,cm}^{-3}$ and $n_\\mathrm{h}\\gtrsim10^{-2}\\,\\mathrm{cm}^{-3}$ close to the MW disk and $P/k_\\mathrm{B}\\lesssim10^2\\ldots10^3\\,\\mathrm{K\\,cm}^{-3}$ and $n_\\mathrm{h}\\sim10^{-4}\\ldots10^{-3}\\,\\mathrm{cm}^{-3}$ at higher Galactic latitudes. Typical masses of the larger structures, e.g., at $(l,b)=(37\\degr,-38\\degr)$, would be few times $10^4\\,M_\\odot$, the smaller clouds could have a few $10^3\\,M_\\odot$.  In their analysis of the clouds at the tip of the MS, \\citet{stanimirovic08} estimate time and length scales for the physical processes important for interacting clouds. While at a first glance our objects have similar properties (i.e., density, size, and line widths), the surrounding medium makes a big difference. Particularly, we observe signs of ongoing ram-pressure interaction (the head--tail structures), which were absent in the MS sample. This is not surprising, given the high densities of the halo material at smaller distances compared to the MS. Hence, in our case the Kelvin--Helmholtz time scales are larger (about 2\\,Gyr), while thermal instabilities in principle should have similar time scales of a few tens of Myr.  \\citet{heitsch09} used simulations to study the stability of neutral clouds against ram-pressure for different scenarios (wind-tunnel and free-fall in the MW potential). While both cases are not exactly applicable for the GCN clouds, some of the authors' basic conclusions might be of interest for us. First, if the initial masses are on the order of $10^3$ to $10^4\\,M_\\odot$, the time scale for complete disruption is on the order of 50 to at most 100\\,Myr (cloud properties in our sample are similar to their simulations and the environmental densities are likely even higher). Although cooling times (and hence the time scale for thermal instabilities) are likely shorter or  comparable, the wind-tunnel simulations show no significant cooling of the clouds. Because the clouds loose material through the ram-pressure interaction, cooling is even less effective. This could explain the small fraction of cold objects in our sample. Note that the ablated material might perhaps form smaller fragments, but these would have much smaller masses than the main cloud. A second important point is that the fragments are most likely not stable (because of the less efficient cooling) and will eventually evaporate. We conclude that ram-pressure interaction can most likely explain the observed cloud properties, though simulations better suited to the GCN scenario would be appreciated.  Under these considerations, complex GCN might be a prime example for warm gas accretion onto the Milky Way, where neutral \\ion{H}{i} clouds or a stream of gas are encountering the Milky Way gas halo. Under certain physical conditions clouds are no longer stable against interaction with the ambient medium, might dissolve and become ionized prior to accretion. This scenario is also suggested by recent simulations \\citep[e.g.,][]{blandhawthorn09}. In that case the GCN objects would just be in a transient state. This would match our observed  $N^\\mathrm{max}_\\ion{H}{i}$--$v_\\mathrm{gsr}$ relation. The lack of \\ion{Na}{i}, which is a tracer for the colder cores of typical clouds, is another sign that the typical life times of the clouds might not be long enough to allow for substantial cooling of the cores --- which explains the under-abundance of cold clouds in our sample compared to most other HVC complexes.   If the hypothesis is true, clouds would not be visible anymore, which immediately raises the question of how much ionized gas mass is contained in complex GCN. Because it is hard to observe, it is still unclear today whether the (global) ionized fraction of accreted material is able to explain the observed constant star-formation rate of the MW. Metal absorption line studies \\citep[using large samples; see][]{benbekhti08,shull11} of the HVC sky will hopefully shed light on this issue. One should also keep in mind that \\ion{Ca}{ii} was detected in sight lines without \\ion{H}{i} emission, hence, there definitely is low-column density material.   We do not only observe a general relation between the upper envelope of the column densities and the radial velocity, but this relation is also significantly found for two subregions (Stripe (a) and N--S). Although the other two subsamples (Region NE and SE) do not individually show a significant relation, all of the slopes are negative, which increases the overall significance. The correlation for Stripe (a) and N--S is surprisingly tight considering the huge scatter in the observed radial velocities in these regions, while for the regions NE and SE, which have much less scatter in radial velocities, the relation is much less constrainted. Furthermore, given the observed (non-)alignments of the HT vectors in each of the regions, we conclude that clouds in the four subsamples may have different origins.   One striking result is the detection of five HT clouds in Stripe~(a) --- the definition of which was based only on the IVC feature near the CND --- pointing eastward. This is compatible with the observed $N^\\mathrm{max}_\\ion{H}{i}$--$v_\\mathrm{gsr}$ relation under the hypothesis that clouds with lower velocities, i.e., which were decelerated for a longer period of time, have already lost more gas through ram pressure. Here, we cannot prove a causal relationship, but it could be that the IVC feature is just the oldest component within Stripe~(a), which has just encountered the MW disk.   The analysis has shown that complex GCN is indeed a very complex object consisting of more than two hundred (mostly individual) tiny clouds, which are often unresolved even with the much better angular resolution of the EBHIS and GASS data compared to any previous study of GCN. We do not observe  diffuse extended 21-cm emission around the isolated clouds, although the column density detection limit of our data would be sufficient to detect this gas phase at reasonable distances. Our observations suggest that the clouds are likely a transient phenomenon in a very turbulent and highly dynamic environment as expected if GCN is really the result of ongoing accretion. It might be that the clouds become ionized prior to accretion, which might render a large fraction of the total gas mass in GCN invisible for the \\ion{H}{i} survey data. The fact that we detect \\ion{Ca}{ii} signatures and other ionized species at velocities consistent with the \\ion{H}{i} structures supports the hypothesis that there is a significant amount of low-column density material around.   The presented work is only the first step toward a deeper understanding of complex GCN as a whole. We are going to use interferometric observations for studying the morphologically interesting cases  at much better angular resolution. This would allow us to obtain better constraints on the true motion of these clouds, because currently there is often an ambiguity between the morphology and observed radial velocity difference between head and tail. It would also be interesting whether unresolved clumps further fragment into individual cores, because this could provide a consistency check for our substructure estimator. Furthermore, deep single-dish observations of some smaller regions could shed light on the question of the missing diffuse extended component. If it turns out that the latter is not physically present, it might be worthwhile to search for additional absorption lines to study the  ionized gas phase to estimate its total mass contribution.  While this work mostly dealt with the general (statistical) properties of our cloud sample, we will next analyze potential subpopulations in more detail (Darmst\\\"{a}dter in prep.). Based on our findings, we suspect that GCN is a mixture of various objects with different origins, which one would need to separate to understand the complex as a whole."
    },
    "1107/1107.2557_arXiv.txt": {
        "abstract": "We investigate the Wightman function, the vacuum expectation values of the field squared and the energy-momentum tensor for a massive scalar field with general curvature coupling in the generalized cosmic string geometry with a compact dimension along its axis. The boundary condition along the compactified dimension is taken in general form with an arbitrary phase. The vacuum expectation values are decomposed into two parts. The first one corresponds to the uncompactified cosmic string geometry and the second one is the correction induced by the compactification. The asymptotic behavior of the vacuum expectation values of the field squared, energy density and stresses are investigated near the string and at large distances. We show that the nontrivial topology due to the cosmic string enhances the vacuum polarization effects induced by the compactness of spatial dimension for both the field squared and the vacuum energy density. A simple formula is given for the part of the integrated topological Casimir energy induced by the planar angle deficit. The results are generalized for a charged scalar field in the presence of a constant gauge field. In this case, the vacuum expectation values are periodic functions of the component of the vector potential along the compact dimension. ",
        "introduction": "Within the framework of grand unified theories, as a result of vacuum symmetry breaking phase transitions, different types of topological defects could be produced in the early universe \\cite{Kibb80,Vile94}. In particular, cosmic strings have attracted considerable attention. Although the recent observational data on the cosmic microwave background have ruled out cosmic strings as the primary source for large scale structure formation, they are still candidates for a variety of interesting physical phenomena such as the generation of gravitational waves \\cite{Damo00}, high energy cosmic rays \\cite{Bhat00}, and gamma ray bursts \\cite{Bere01}. Recently the cosmic strings attracted a renewed interest partly because a variant of their formation mechanism is proposed in the framework of brane inflation~\\cite% {Sara02}.  In the simplest theoretical model describing the infinite straight cosmic string the spacetime is locally flat except on the string where it has a delta shaped curvature tensor. The corresponding nontrivial topology raises a number of interesting physical effects. One of these concerns the effect of a string on the properties of quantum vacuum. Explicit calculations for vacuum polarization effects in the vicinity of a string have been done for various fields \\cite{Hell86}-\\cite{BezeKh06}. Vacuum polarization effects by the cosmic string carrying a magnetic flux are considered in Ref. \\cite% {Guim94}. Another type of topological quantum effects appears in models with compact spatial dimensions. The presence of compact dimensions is a key feature of most high energy theories of fundamental physics, including supergravity and superstring theories. An interesting application of the field theoretical models with compact dimensions recently appeared in nanophysics. The long-wavelength description of the electronic states in graphene can be formulated in terms of the Dirac-like theory in three-dimensional spacetime with the Fermi velocity playing the role of speed of light (see, e.g., \\cite{Cast09}). Single-walled carbon nanotubes are generated by rolling up a graphene sheet to form a cylinder and the background spacetime for the corresponding Dirac-like theory has topology $% R^{2}\\times S^{1}$. The compactification of spatial dimensions serves to alter vacuum fluctuations of a quantum field and leads to the Casimir-type contributions in the vacuum expectation values of physical observables (see Refs. \\cite{Most97,Eliz94} for the topological Casimir effect and its role in cosmology). In the Kaluza-Klein-type models, the topological Casimir effect induced by the compactification has been used as a stabilization mechanism for the size of extra dimensions. The Casimir energy can also serve as a model for dark energy needed for the explanation of the present accelerated expansion of the universe. The influence of extra compact dimensions on the Casimir effect in the classical configuration of two parallel plates has been recently discussed for the case of a scalar field \\cite{Chen06}, for the electromagnetic field with perfectly conducting boundary conditions \\cite{Popp04}, and for a fermionic field with bag boundary conditions \\cite{Bell09}.  In this paper we shall study the configuration with both types of sources for the vacuum polarization, namely, a generalized cosmic string spacetime with a compact spatial dimension along its axis (for combined effects of topology and boundaries on the quantum vacuum in the geometry of a cosmic string see \\cite{Brev95,Beze06,Beze07}). For a massive scalar field with an arbitrary curvature coupling parameter, we evaluate the Wightman function, the vacuum expectation values of the field squared and the energy-momentum tensor. These expectation values are among the most important quantities characterizing the vacuum state. Though the corresponding operators are local, due to the global nature of the vacuum, the vacuum expectation values carry an important information about the the global properties of the bulk. In addition, the vacuum expectation value of the energy-momentum tensor acts as the source of gravity in the semiclassical Einstein equations. It therefore plays an important role in modelling a self-consistent dynamics involving the gravitational field. The problem under consideration is also of separate interest as an example with two different kind of topological quantum effects, where all calculations can be performed in a closed form.  We have organized the paper as follows. The next section is devoted to the evaluation of the Wightman function for a massive scalar field in a generalized cosmic string spacetime with a compact dimension. Quasiperiodic boundary condition with an arbitrary phase is assumed along the compact dimension . By using the formula for the Wightman function, in section \\ref% {Sec:phi2} we evaluate the vacuum expectation value of the field squared. This expectation value is decomposed into two parts: the first one corresponding to the geometry of a cosmic string without compactification and the second one being induced by the compactification. The vacuum expectation value of the energy-momentum tensor is discussed in section \\ref% {Sec:EMT}. Section \\ref{Sec:VacEn} is devoted to the investigation of the part in topological vacuum energy induced by the planar angle deficit. Finally, the results are summarized and discussed in section \\ref{sec:Conc}. ",
        "conclusions": "\\label{sec:Conc}  In the present paper we have investigated the one-loop quantum effects for a massive scalar field with general curvature coupling parameter, induced by the compactification of spatial dimensions in a generalized cosmic string spacetime. It is assumed that along the compact dimension the field obeys quasiperiodicity condition with an arbitrary phase. As the first step for the investigation of vacuum densities we evaluate the positive frequency Wightman function. This function gives comprehensive insight into vacuum fluctuations and determines the response of a particle detector of the Unruh-DeWitt type in a given state of motion. For a massive field and for general value of the planar angle deficit, the Wightman function is given by formula (\\ref{WF4}). The $l=0$ term in this expression corresponds to the Wightman function in the geometry of a cosmic string without compactification and, hence, the topological part is explicitly extracted. As the compactification under consideration does not change the local geometry, in this way the renormalization for the VEVs in the coincidence limit is reduced to that for the standard cosmic string geometry without compactification. For integer values of the parameter $q$, the Wightman function is expressed as an image sum of the corresponding function in the Minkowski spacetime with a compact dimension and is given by Eq. (\\ref{WF4Sp}% ).  The VEV of the field squared is decomposed as the sum of the pure string part and the correction due to the compactification. For a massive field the string part is given by (\\ref{phi2s}). Since the geometry is locally flat, this part does not depend on the curvature coupling parameter. It is positive for $q>1$ and diverges on the string like $1/r^{D-1}$. This divergence may be regularized considering a more realistic model of the string with nontrivial inner structure. The topological part in the VEV\\ of the field squared is given by the expression (\\ref{phi2t}) and it is finite everywhere including the points on the string. In dependence of the value of the phase $\\beta $ in the quasiperiodicity condition, this part can be either positive or negative. In particular, the topological part is positive for an untwisted scalar and it is negative for a twisted scalar. At distances from the string larger than the length of the compact dimension, the topological part in the VEV of the field squared approaches to the corresponding quantity in the Minkowski spacetime with a compact dimension. For points near the string we have the simple asymptotic relation $\\langle \\varphi ^{2}\\rangle _{\\text{t}}\\approx q\\langle \\varphi ^{2}\\rangle _{\\text{t}}^{% \\text{(M)}}$.  The VEV of the energy-momentum tensor is investigated in section \\ref% {Sec:EMT}. This VEV is diagonal and, similar to the case of the field squared, it is decomposed into the pure string and topological parts, given by expressions (\\ref{EMTs}) and (\\ref{EMTt}), respectively. We have explicitly checked that the topological part satisfies the covariant conservation equation and its trace is related to the VEV of the field squared by the standard formula. For a massive field and for general value of the parameter $q$, we give a closed expression for the pure string part in the VEV of the energy-momentum tensor in arbitrary number of dimensions. The latter generalizes various special cases previously discussed in the literature. At large distances from the string, the topological part coincides with the corresponding result in the Minkowski spacetime and it dominates in the total VEV. In this limit the VEV of the energy-momentum tensor does not depend on the curvature coupling parameter and the corresponding energy density is negative/positive for untwisted/twisted scalar fields. The topological part is finite on the string and for points near the string the leading term in the corresponding asymptotic expansion is given by (\\ref{EMTtaxis}). The pure string part in the VEV of the energy density diverges on the string as $1/r^{D+1}$ and near the string it dominates in the total VEV. For a conformally coupled scalar field the corresponding energy density is negative, whereas for a minimally coupled field the energy density is positive for small values of the parameter $q$ and becomes negative for large values of $q$. The numerical examples are given for the simplest Kaluza-Klein-type model with a single extra dimension. They show that he nontrivial topology due to the cosmic string enhances the vacuum polarization effects induced by the compactness of spatial dimensions for both the field squared and the vacuum energy density. For a charged scalar field, in the presence of a constant gauge field the expression for the topological parts are obtained from the formulas given above by the replacement $\\beta \\rightarrow \\beta ^{\\prime }$ with $\\beta ^{\\prime }$ defined by (\\ref{beta}). In this case the topological parts are periodic functions of the component of the gauge field along the compact dimension. This is an analog of the Aharonov-Bohm effect.  As a result of the non-integrable divergence of the energy density in the pure string part on the string, the corresponding total vacuum energy is divergent. In section \\ \\ref{Sec:VacEn} we give a closed expression, Eq. (% \\ref{Es0}), for the vacuum energy in the region $r\\geqslant r_{0}>0$. The topological part in the VEV of the energy-momentum tensor can be further decomposed into the Minkowskian and string induced parts. The latter is finite on the string and vanishes at large distances from the string. As a result, the total vacuum energy corresponding to this part is finite. This energy is given by the expressions (\\ref{EST2}) and (\\ref{ESTm0}) for massive and massless fields, respectively. In these expressions the dependence on the parameter $q$ is simply factorized. The string induced part in the topological energy does not depend on the curvature coupling parameter and it is negative/positive for untwisted/twisted scalars.  In a way similar to that described above, one can consider the topological Casimir effect for a cosmic string in de Sitter spacetime with compact spatial dimensions. The vacuum polarization effects induced by the presence of the string in uncompactified de Sitter spacetime have been recently discussed in Ref. \\cite{Beze09}. The topological Casimir densities in de Sitter spacetime with toroidally compactified spatial dimensions are investigated in Ref. \\cite{Saha08}. It has been shown that the curvature of the background spacetime decisively influences the behavior of the topological parts in the VEVs of the field squared and the energy density for lengths of compact dimensions larger than the curvature scale of the spacetime.  In this paper the string geometry is taken as a static, given classical background for quantum matter fields. This approach follows the main part of the papers where the influence of the string on quantum matter is investigated (see references \\cite{Hell86}-\\cite{Guim94}). Of course, in a more complete approach the dynamics of the cosmic string should be taken into account. In the simplest model, the cosmic string dynamics can be described by the Nambu action (see, for instance, \\cite{Vile94}). If the scalar field under consideration interacts with the Higgs field inside the string core, then, within this model, the total action will contain also the term describing the interaction of the scalar field with the vibrational modes of the string. This would be the further development of the model under discussion. The results obtained in the present paper are the first step to this more general problem. Another development would be the investigation of the back-reaction effects of the quantum energy-momentum tensor on the gravitational field of the cosmic string. For the geometry of infinitely thin straight cosmic string the back-reaction for conformal fields has been discussed in \\cite{Iell97,Hisc87} by using the linearized semiclassical Einstein equations. It would also be interesting to generalize the vacuum polarization calculations of the present paper for the models with nontrivial string core. For a general cylindrically symmetric static model of the string core with finite support this can be done in a way similar to that used in \\cite{Beze06} for the geometry of a straight cosmic string."
    },
    "1107/1107.2885_arXiv.txt": {
        "abstract": "We present the results of an extensive study of the detectability of Earth-sized planets and super-Earths in the habitable zones of cool and low-mass stars using transit timing variation method. We have considered a system consisting of a star, a transiting giant planet, and a terrestrial-class perturber, and calculated TTVs for different values of the parameters of the system. To identify ranges of the parameters for which these variations would be detectable by {\\it Kepler}, we considered the analysis presented by Ford et al. (2011, ArXiv:1102.0544) and assumed that a peak-to-peak variation of 20 seconds would be within the range of the photometric sensitivity of this telescope. We carried out simulations for resonant and non-resonant orbits, and identified ranges of the semimajor axes and eccentricities of the transiting and perturbing bodies for which an Earth-sized planet or a super-Earth in the habitable zone of a low-mass star would produce such TTVs. Results of our simulations indicate that in general, outer perturbers near first- and second-order resonances show a higher prospect for detection. Inner perturbers are potentially detectable only when near 1:2 and 1:3 mean-motion resonances. For a typical M star with a Jupiter-mass transiting planet, for instance, an Earth-mass perturber in the habitable zone can produce detectable TTVs when the orbit of the transiting planet is between 15 and 80 days. We present the details of our simulations and discuss the implication of the results for the detection of terrestrial planets around different low-mass stars. ",
        "introduction": "\\label{intro}   Cool and low-mass stars (such as M dwarfs) present the most promising candidates for searching for Earth-like planets and super-Earths. These stars show large reflex accelerations in response to an orbiting body, and their low surface temperatures, which place their liquid water habitable zones (Kasting et al 1993) at close distances of 0.1--0.2 AU (corresponding to orbital periods of 20--50 days), cause their light-curves to show large decrease when they are transited by a close-in planet. In the past few years, these advantages have resulted in the detection of several super-Earths around low-mass stars including a 6.0--7.5 Earth-mass planet around the star GJ 876 (Rivera et al. 2005, 2010; Correia 2010), 4--6 planets with masses ranging from 1.7 to 7.0 Earth-masses around star GL 581 (Bonfils et al 2005; Udry et al 2007; Mayor et al 2009; Vogt et al 2010), and two transiting planets around M stars GJ 436 (Butler et al 2004; Gillon 2007) and GJ 1214 (Charbonneau 2009).  The existence of systems such as GJ 876 and GL 581, where an M star is host to multiple planets in short-period orbits, suggests that other single-planet low-mass stars may also host additional objects. In systems where one (or more) of these bodies transits, the interaction between this object and other planets of the system results in the exchange of energy and angular momentum, and creates small, short-term oscillations in their semimajor axes and orbital eccentricities. On the transiting planet, these oscillations manifest themselves as variations in the time and duration of the planet's transit. In this paper, we discuss the possibility of the detection of these transit timing variations when they are generated by an Earth-like planet or a super-Earth in the habitable zone of a low-mass star.  Measurements of transit timing variations can be used in detecting small, non-transiting planets (Miralda-Escud\\'e 2002, Schneider 2003, Holman and Murray 2005, Agol et al 2005, Heyl and Gladman 2007). Despite the complication in extracting the mass and orbital elements of the perturbing body from the values of TTVs (Nesvorn\\'y and Morbidelli 2008, Nesvorn\\'y 2009, Nesvorn\\'y and Beaug\\'e 2010, Meschiari and Laughlin 2010, Veras et al 2011), this method has also been recognized as a potential mechanism for detecting Trojan planets (Ford and Gaudi 2006, Ford and Holman 2007, Haghighipour et al, in preparation), as well as the moons of transiting gas-giants (Sartoretti and Schneider 1999, Simon et al 2007, Kipping 2009a\\&b). TTVs are also sensitive to the precession (Miralda-Escud\\'e 2002) and inclination of the orbit of the transiting object (Payne et al 2010), and in close, eclipsing binary systems, they can be used to detect circumbinary planets (Schneider and Chevreton 1990, Schneider and Doyle 1995, Doyle et al 1998, Doyle and Deeg 2003, Deeg et al 2008, Sybilski et al 2010, Schwarz et al 2011).  In this paper, we focus our study on the detection of TTVs using {\\it Kepler} space telescope. {\\it Kepler} monitors more than 150,000 stars with spectral types of A to M, and with apparent magnitudes brighter than 14 (Oshagh et al 2010; Batalha et al 2010, Coughlin et al 2011). A survey of the discovery-space of this telescope indicates that low-mass stars such as M dwarfs, constitute a significant number of its targets. The large number of the targets of this telescope, combined with its unprecedented photometric precision, and the long duration of its operation, makes {\\it Kepler} ideal for searching for transit timing variations. A recent success, for instance, has come in the form of detecting variations in the orbital periods of transiting planets around stars Kepler-9 (Holman et al 2010) and Kepler-11 (Lissauer et al 2011). Further analysis of the currently available data from this telescope also points to the existence of many additional candidates that have potentially detectable variations in the times of their transits (Steffen et al 2010,  Ford et al 2011).  As shown by Agol et al (2005), Steffen and Agol (2005), and  Agol and Steffen (2007), the amplitude of TTV is greatly enhanced when the transiting and perturbing planets are near a mean-motion resonance (MMR). In this case, transit times vary as a function of the ratio of the masses of the two planets which enhances the sensitivity of the TTV method by one to two orders of magnitude. For instance, as shown by Agol et al (2005), around a solar-mass star, an Earth-mass planet can create TTVs with amplitudes of up to 180 s when near an exterior 2:1 MMR with a transiting Jupiter-size body in a 3-day orbit. In the recently detected Kepler-9 system where two Saturn-sized planets are near a 2:1 resonance, the rate of the transit timing variation on the inner planet was even higher and reached to approximately 50 minutes. This is two orders of magnitude higher than the $\\sim$10 s transit timing precisions that was achieved by a 1.2  m ground-based telescope in the measurements of the transit timing of HD 209458 (Brown et al 2001). The fact that a perturber near a mean-motion resonance can produce large TTVs implies that such resonant systems will have a higher chance of being detected by the TTV method. In this paper, we will pay especial attention to such resonant orbits.   The outline of this paper is as follows. In section 2, we present the details of our methodology for calculating TTVs. Since we are interested in the detectability of an Earth-like planet or a super-Earth in the habitable zone of a star, we will analyze the dependence of the magnitudes of TTVs on different parameters of a system and present the results in section 3. Section 4 has to do with the implications of this analysis for the detection of planets in the habitable zone, and section 5 concludes this study by presenting a summary and reviewing the results. ",
        "conclusions": "We carried out an extensive study of the possibility of the detection of small planets in the HZ of low-mass stars with {\\it Kepler} space telescope and transit timing variation method. With its high precision photometry, {\\it Kepler} has the capability of detecting small TTVs such as those created by terrestrial-class perturbers. We assumed that TTVs with amplitudes as small as  20 s can be detected by {\\it Kepler}, and identified regions of the parameter space where the variations in the time of the transit of a giant planet in a short-period orbit around a low-mass star would be larger than this value.  In general, the lowest detectable amplitude of TTVs depends upon several quantities including the magnitude of the central star, the depth of the transit, the durations of ingress and egress, and the number of transits observed. There will also be systematics that will affect the sensitivity of the telescopes at different times. The 20 s detectability limit considered in this study corresponds to the best transit timing precision obtained by Ford et al (2011) in their statistical analysis of the transit timing of {\\it Kepler}'s planetary candidates. Analyzing the data obtained from 1235 potential transiting planets (Borucki et al 2011) during the first four months of the operation of this telescope, Ford et al obtained a transit timing accuracy ranging from approximately 20 s to 100 minutes (figure 1 of Ford et al 2011). The 20 s precision is also consistent with the analysis of the light curves of Kepler-4 through Kepler-8 as given by Kipping and Bakos (2011). As shown by these authors, the timing precision for the transiting planets in these systems ranges between 20 and 250 seconds. It is worth noting that systems with TTVs larger than 15-20 s may also be detected by CoRoT. This telescope has been monitoring more than 12000 dwarf stars with apparent magnitudes ranging from 11 to 16.  Although many non-resonant systems produced detectable TTVs, the results of our study indicate that as expected, perturbers near mean-motion resonances with the transiting planet show a higher prospect of detection. This can also be seen in the system of Kepler-9 where the two TTV-detected planets of this system are near a 2:1 MMR (Holman et al 2010). Our study shows that interior perturbers with masses ranging from 1 to 10 $M_\\oplus$ can produce detectable TTVs when near 1:2 and 1:3 MMRs. However, in the majority of our systems, these interior perturbers were not in the HZ. When in an outer orbit, these planets can be in the HZ and produce detectable TTVs when in first- and second-order resonances. For instance, for a typical M star with a mass of approximately 1/3 of the Sun where the HZ extends between 0.1 and 0.2 AU, perturbers in 3:2, 2:1, 5:3, and 3:1 MMRs produced large-magnitude TTVs when in the system's HZ. For the systems studied here, a perturber with a mass ranging from 1 to 10 Earth-masses can be detected in the HZ when the transiting planet is in 15-80 day orbits.  In this study we considered the planetary system to be planar and consists of only three bodies. These assumptions may limit the applicability of our results. For instance, an inclined perturber will cause the inclination of the orbit of the transiting planet to vary which in turn affects the variations in the times of its transits (Payne et al 2010). Also, as suggested by the simulations of planet formation around low-mass stars, these stars may be hosts to more than two planets (e.g., GJ 876 hosts four planets, and GL 581 has 4--6 super-Earths). Depending on their masses and orbital elements, additional bodies may play a role in the magnitude of TTVs.  Another assumption considered in our study is the existence of a giant planet in a short-period orbit around a low-mass star. In general, the low masses of circumstellar disks around these stars, combined with the short lifetime of the gas in such disks, casts doubt in the possibility that giant planets may exist in these systems and in close-in orbits (Laughlin et al 2004). The fact that giant planets have been observed around M stars [e.g., GJ 876 with two Jupiter-like planets and a Uranus-mass body in approximately 30, 60, and 120 days orbits (Rivera et al. 2005, 2010), or HIP 57050 with a Saturn-mass planet in a 42 days orbit (Haghighipour et al 2010)] suggests that these planets might have formed at the outer regions of the disk where more material was available, and migrated to their current orbits. As shown by Zhou et al (2005), Fog and Nelson (2005, 2006, 2007a \\& b, 2009), Raymond et al (2006), Mandell et al (2007), and Kennedy and Kenyon (2008), terrestrial planets may form during the migration of a giant planet around a Sun-like star. These planets may be captured in resonance with the giant planet and migrate to smaller orbits. Recently Haghighipour and Rastegar (2011) carried out similar simulations around M stars. By allowing the giant planet to migrate to orbits with periods of 3-5 days, these authors showed that around low-mass stars, resonance capture during the migration of the giant planet seems to be the most viable mechanism for the formation of a system that consists of a transiting giant planet and a small interior perturber. In case of an outer perturber, although it is possible for small planets to migrate and capture larger bodies in resonance (e.g., GJ 876 where the outermost planet is a Uranus-mass object in a 4:2:1 resonance with the two inner Jupiter-size planets), it is more probable that the smaller, outer perturber is accompanied by another larger planet in an exterior mean-motion resonance. How such a second exterior perturber affects the results presented in this paper is the subject of our currently on-going research."
    },
    "1107/1107.3621_arXiv.txt": {
        "abstract": "Systematically studying all the {\\it RXTE}/PCA observations for GRS 1915+105 before November 2010, we have discovered three additional patterns in the relation between Quasi-Periodic Oscillation (QPO) frequency and photon energy, extending earlier outcomes reported by \\citet{Qu10}. We have confirmed that as QPO frequency increases, the relation evolves from the negative correlation to positive one. The newly discovered patterns provide new constraints on the QPO models. ",
        "introduction": "GRS 1915+105, discovered by WATCH instrument on board {\\emph{GRANAT}} in 1992 \\citep{Castro92} and located in our galaxy at an estimated distance of $9\\pm3$ kpc \\citep{Chapuis04}, is a low-mass X-ray binary containing a spinning accreting black hole \\citep{Zhang97} of mass about $14\\pm4$ M$_{\\odot}$ and a K-M III giant star of mass $0.8\\pm0.5$ M$_{\\odot}$ as the donor \\citep{Harlaftis04, Greiner01a}. The orbital separation and period of this binary are, respectively, about $108\\pm4$ R$_{\\odot}$ and 33.5 days \\citep{Greiner01b}. Serving as a famous microquasar, GRS 1915+105 produces superluminal radio jets \\citep{Mirabel94, Fender99}. It shows various X-ray light curves and complex timing phenomena. Based on the appearance of light curves and color-color diagrams, the behaviors of GRS 1915+105 can be classified into 12 classes. The variability of the source can be further reduced to transitions between three basic states (A, B, and C) \\citep{Belloni00}. Of these 12 classes, class $\\chi$ (state) is most commonly observed \\citep{Belloni00}. It shows characteristics exclusively of state C, the state which is steady in the X-rays and lies in a rather hard part of the color-color diagram. It is the state when the low-frequency ($\\sim0.5-10$ Hz) QPOs (LFQPOs) are most frequently observed \\citep[e.g.,][]{Muno99}, providing an idea site for studying LFQPOs.  Much effort has been made for exploring the origins of the LFQPOs of GRS 1915+105. It was found that the QPO frequency was positively correlated with the fluxes of the individual components and spectral total flux \\citep[e.g.,][]{Chen97, Markwardt99, Muno99, Trudolyubov99, Reig00, Tomsick01}. \\citet{Muno01} confirmed that QPO frequency was tightly correlated with the source flux, and, however, they also found that for some observations the QPO frequency was not correlated with the flux. It was found that the QPO amplitude was inversely correlated with the source flux or QPO frequency \\citep[e.g.,][]{Muno99, Reig00, Trudolyubov99}. \\citet{Muno99} and \\citet{Rodriguez02a} reported that as QPO frequency increased, the temperature of the inner accretion disk increased and the radius of the inner accretion disk decreased. These results indicate that the LFQPO is linked to both the accretion disk and the region where the power law component is produced. However, most of these results are related to models. As a model-independent means, it is meaningful to study the correlations between photon energy and QPO parameters including its amplitude and frequency. Some authors found that the QPO amplitude increased with photon energy and it turned over in high energy bands in some cases \\citep[e.g.,][]{Tomsick01, Rodriguez02b, Rodriguez04, Zdziarski05}. \\citet{Qu10} studied the LFQPOs of GRS 1915+105 in class $\\chi$ state \\citep{Belloni00} and found that as the centroid frequency of QPO increased the correlation between QPO frequency and photon energy evolved from a negative correlation to a positive one.  Nevertheless, systematic studies on the energy dependence of the LFQPO frequency in GRS 1915+105 have never been done. In this work, using all the data of {\\it RXTE}/PCA of GRS 1915+105 before November 2010, we have investigated the correlation between photon energy and QPO frequency throughout. The data reduction methods are described in \\S 2, the results are presented in \\S 3, while a simple discussion and the conclusion are given in \\S 4. ",
        "conclusions": "By systematically analyzing all the {\\it RXTE}/PCA observations of GRS 1915+105 before November 2010, we have found that there are five typical patterns of the relation between QPO frequency and photon energy: the negative correlation (P1) and the positive one (P5), as well as three combined relations of them (P2, P3 and P4). The P1 and P5 were firstly found by \\citet{Qu10} and the intermediate patterns, P2, P3 and P4, were newly discovered in this work. Besides, we have confirmed in large sample the result reported by \\citet{Qu10}: with increasing of the centroid frequency of QPO, the relation between QPO frequency and photon energy evolves from P1 to P5.  Following \\citet{Qu10}, we apply several models to the detected patterns between the QPO frequency and photon energy. \\citet{Titarchuk00} proposed the global disk model (GDM) to interpret the LFQPOs in X-ray binaries, in which they argued that the disk oscillations were the result of gravitational interaction between the the central compact objects and the disk and the QPO frequency was determined by the mass of the central object. Thus, the QPO frequency determined by the GDM is expected to be independent of the photon energy and therefore, this model can explain those QPOs whose coefficients are zero, as demonstrated in bottom panel of Figure \\ref{fig:fre_k}. The radial and orbital oscillation model (ROOM) is a typical model for QPOs \\citep{Nowak93, Nowak94}. According to this model, the oscillation frequency will vary with the the disk radius or temperature, resulting in that the QPO frequency will vary with the photon energy, because photons with different energies are from different radii. Simply assuming the disk oscillation frequency to be the Keplerian frequency at the inner disk radius, it is natural to explain the relation P5: with decreasing of the inner disk radius, the QPO frequency increases, and meanwhile the photon energy increases, because the temperature at inner disk edge also increases. Fortunately, \\citet{Muno99} reported that in this source the QPO parameters were indeed correlated with the disk parameters. Other models such as the drift blob model \\citep[DBM;][]{Bottcher98, Bottcher99} can also be used to explain the energy dependence of QPO frequency. In the DBM, the blobs drifting inward through an inhomogeneous hot inner disk or corona would cause the QPO frequency to increase with energy. Both the ROOM and DBM can explain the positive correlation between the QPO frequency and photon energy, but they are frustrated by the negative correlation (P1) and other complicated correlations (P2, P3, and P4). There is no model for interpreting all the relations, which perhaps indicates that the QPOs with different relations between the QPO frequency and photon energy result from different mechanisms.  This work presents various patterns of relation between the QPO frequency and photon energy in the famous microquasar GRS 1915+105, indicating that the mechanisms for QPOs are complicated. On the other hands, we extend the investigation on the relations made by \\citet{Qu10} and the newly discovered relations provide new clues on QPO models."
    },
    "1107/1107.4371_arXiv.txt": {
        "abstract": "In this work we analyze the physical properties of a sample of 153 star forming galaxies at z$\\sim$0.84, selected by their H$\\alpha$ flux with a NB filter. B-band luminosities of the objects are higher than those of local star forming galaxies. Most of the galaxies are located in the blue cloud, though some objects are detected in the green valley and in the red sequence. After the extinction correction is applied virtually all these red galaxies move to the blue sequence, unveiling their dusty nature. A check on the extinction law reveals that the typical extinction law for local starbursts is well suited for our sample but with E(B-V)$_{stars}$=0.55 E(B-V)$_{gas}$. We compare star formation rates (SFR) measured with different tracers (H$\\alpha$, FUV and IR) finding that they agree within a factor of three after extinction correction. We find a correlation between the ratios SFR$_{FUV}$/SFR$_{H\\alpha}$, SFR$_{IR}$/SFR$_{H\\alpha}$ and the EW(H$\\alpha$) (i.e. weighted age) which accounts for part of the scatter. We obtain stellar mass estimations fitting templates to multi-wavelength photometry. The typical stellar mass of a galaxy within our sample is $\\sim$10$^{10}$M$_{\\odot}$. The SFR is correlated with stellar mass and the specific star formation rate (sSFR) decreases with it, indicating that massive galaxies are less affected by star formation processes than less massive ones. This result is consistent with the {\\em downsizing} scenario. To quantify this {\\em downsizing} we estimated the {\\em quenching} mass M$_{Q}$ for our sample at z$\\sim$0.84, finding that it declines from M$_{Q}\\sim$10$^{12}$M$_{\\odot}$ at z$\\sim$0.84 to M$_{Q}\\sim$8$\\times$10$^{10}$ M$_{\\odot}$ at the local Universe. ",
        "introduction": "During the last two decades the cosmic Star Formation History of the Universe has been widely studied in order to better constraint galaxy formation and evolution models. Large-area surveys as well as the use of larger telescopes have consolidated our knowledge at low-intermediate redshifts (z=0.0-1.0) \\citep[see][]{Hop06}. Several measurements exist at higher redshifts \\citep{PG05,Geach08, Reddy09, Hayes10} and at present we have started to probe the most distant Universe at z$\\sim$7--8 \\citep{Bouwens09,Bouwens10arXiv}.  In general, the Star Formation Rate density (SFRd) history from the Local Universe to z$\\sim$1 is well accepted. Near z$\\sim$1, where the rise in SFRd from the local Universe slows down, there are several measurements obtained through samples selected in a variety of ways and using different Star Formation Rate (SFR) tracers. Results from H$\\alpha$, UV and IR agree reasonably, although with higher dispersion than at lower redshifts \\citep{Garn10} . This scattering is originated (at least partially)  due to the two aforementioned processes: the sample selection or/and the SFR estimation. The estimations of the SFRd at this redshift have been measured mainly through UV \\citep{Lilly96,Connolly97,Cowie99,Wilson02,Schimi05}, IR \\citep{Flores99,Lefloch05,PG05} and H$\\alpha$ \\citep{gla99,Yan99,Hop00,Tresse02,Do06,Villar08,Sobral09, Ly11}. Therefore, it is important to constraint the potential differences that arise when one or another tracer is used to estimate SFRs.  Estimations of SFR through the UV flux are very sensitive to the extinction correction. The UV also probe older populations than H$\\alpha$, thus being more sensitive to recent star formation history \\citep{Cal05}. The IR on the other hand is not affected by extinction but is very model dependent \\citep{Barro11,Barro11b}. Moreover, it is not well understood how the old population of stars contributes to the IR emission, but it could be a significant fraction \\citep{DaCunha08,Salim09}. The H$\\alpha$ line is one of the best estimators as it is sensitive only to very young stars, not affected by recent star formation history, which may still be detected by UV or IR. It has the problem that at z$>$0.5 it moves to the near Infrared (nIR) domain, where large amounts of spectroscopy are still difficult to obtain, though several nIR multi-object spectrographs for 8-10 meter class telescopes are coming in the next years. In addition, different methods have been used for measuring the total H$\\alpha$ line flux, which could lead to discrepancies if some effects are not properly corrected. On one hand we have spectroscopy, long- or multi- slit or through fibers, where aperture corrections are needed to recover the total flux \\citep[see for example][]{Do06,Erb06}. On the other hand we have the slitless spectroscopy \\citep{Yan99,Hop00} and narrow band \\citep{Villar08,Sobral09,Ly11} techniques, which have the advantage that the total flux of the object is recovered and no aperture corrections are needed. Although it is needed to correct from the Nitrogen contribution if the filter is not narrow enough, this effect is easier to estimate than aperture corrections. Obviously, the extinction is also important at this wavelength, though less than in the UV.  Thus, the H$\\alpha$ estimator is the best suited to study the instantaneous star formation. As mentioned before the problem is that at this redshift the line is observed in the nIR and little data is available up to date. Then, it is also interesting to asses if UV and IR provide SFRs comparable to those of H$\\alpha$, given the large amount of data in these wavelengths available today.  The selection methods of the samples are also different and target different populations. Selecting the sample in the UV for example, implies a bias against very obscured galaxies, which may not be detected unless very deep observations are carried out. On the other hand, the IR selected samples will favor the selection of objects with large amounts of dust, missing the blue and dust-free objects. Selections based in the H$\\alpha$ line will select the objects with ongoing star formation and thus only star forming galaxies (and AGN) will be selected.  Thus, to study the population of galaxies that are actively forming stars it is necessary to have a well defined sample of star forming galaxies and one of the best technique today is the use of narrow band filters targeting H$\\alpha$.  A population of star forming galaxies selected in this way is ideal to study the SFR sequence, which refers to the correlation that exists between SFR and stellar mass. This correlation has been found at a wide range of redshifts although there has been found evolution with respect to this parameter \\citep[see][for a review]{Dutton10}. While the slope is almost constant, the SFR zeropoint increase from the local Universe to redshift z$\\sim$2. From this redshift on, the trend remains almost constant with little evolution. There exist some discrepancy in the slope, which is somewhat lower than unity in all cases \\citep[see for example][]{Noeske07a,Elbaz07,Salim07}.  However, some works do not find this correlation. \\cite{Caputi06} do not find any trend for their MIPS selected sample at z$\\sim$2. More recently, \\cite{Sobral10} did not find any evidence of this sequence for their HiZELS sample at z$\\sim$0.84, selected with a narrow-band filter targeting H$\\alpha$.  This correlation implies that the slope of specific SFR (sSFR) versus mass is higher than -1. Obviously, if no correlation between stellar mass and SFR is found the slope of sSFR vs. stellar mass is simply -1. The slope of this correlation is very important, as it tells us how decreases the importance of star formation over the already formed stellar mass.   In this work we use a narrow band selected sample of star forming galaxies at z$\\sim$0.84, presented in \\cite{Villar08}, to compare SFRs obtained from different tracers and to study the relation between stellar mass and SFR. The sample is very well suited for this study as: i) it is directly selected by star formation, so we are not biasing the population of star forming galaxies; ii) the use of the narrow band filter technique provides us reliable H$\\alpha$ SFRs to compare with FUV and IR estimations.  This paper is structured as follows. In Section~\\ref{sec:sample}, we present the sample and the available datasets. In Section~\\ref{sec:agn}, we describe the methods used to get rid of AGN contaminants. Absolute magnitudes and color are presented in Section~\\ref{sec:photo}. Section ~\\ref{sec:sfr} presents a comparison of SFRs obtained through different estimators as well as a check on the extinction law more suited to our sample. Section~\\ref{sec:mass} presents the stellar masses and their relation with star formation rate. Finally, we summarize our results and conclusions in Section~\\ref{sec:conclusions}.  Throughout this paper we use AB magnitudes. We adopt the cosmology {\\em H$_{0}$} = 70 km s$^{-1}$ Mpc$^{\u22121}$, $\\Omega _{m}$ = 0.3 and $\\Omega _{\\lambda}$ = 0.7. ",
        "conclusions": "\\label{sec:conclusions}  In this work we have analyzed the properties of an H$\\alpha$ selected sample of star-forming galaxies at z$\\sim$0.84, focusing on the star formation and stellar mass.  We have discarded the AGN contaminants through two criteria: X-ray luminosities and IRAC colors. We find seven counterparts in X-ray, though three of them present very low fluxes, compatible with being originated by star formation processes. Thus, we only discard the four objects with fluxes high enough to have an AGN origin. Using IRAC colors we find another ten objects (one of them already detected in X-rays) that fulfill the criterion to be considered AGN. A total of thirteen objects are finally discarded.  The objects of our sample present a median $M_{B}$=-20.5$\\pm$0.9$^{m}$, brighter by more than one magnitude than the UCM local sample of star forming galaxies. Most of the galaxies belong to the blue sequence with a small fraction of objects in the green valley and the red sequence. Once the extinction corrections are applied all these red objects except two move to the blue sequence, unveiling their dusty nature.  A check on the extinction law reveals that the \\cite{Cal00} extinction law is appropriate for our objects, but with E(B-V)$_{stars}$=0.55 E(B-V)$_{gas}$.  The H$\\alpha$ SFR, without applying the extinction correction, presents values in the range 2--10 M$_{\\odot}$. We have also estimated SFRs with FUV and IR. In the first case, the non-extinction corrected FUV underestimates the SFR with respect to H$\\alpha$. The opposite case is given for the IR, which overestimates the SFR with respect to H$\\alpha$. These discrepancies are mainly driven by the extinction. Once we apply the extinction correction to both FUV and H$\\alpha$ estimations, all SFR tracers agree within a factor of three and the highest SFRs reach several hundreds solar masses per year.  The scattering between the different tracers are correlated with the H$\\alpha$ equivalent width. This  can be explained through the different weighted age of the objects (which is related to the EW) and the fact that FUV an IR SFRs are sensitive to a longer time range than H$\\alpha$, being more affected by older populations.  We have estimated stellar masses for our sample, finding that the median value is M$_{\\star}$=1.4$^{+4.6}_{-1.1} \\times 10^{10} $M$_{\\odot}$, in good agreement with the result obtained by S10 for a sample selected with similar criteria as ours. The typical mass found by \\cite{PG08} for an IRAC selected sample at this redshift is ten times higher than this value.  Our sample shows a trend between SFR and stellar mass. The slope of this trend is in good agreement with the value obtained by \\cite{Noeske07b} and is flatter than the \\cite{Elbaz07} result. The trend is very similar to that of the local Universe, although shifted to higher values of SFR. This indicates that, for the same stellar mass M$_{*}$, star forming galaxies at z$\\sim$0.84 are under stronger star formation episodes than their local analogous.  The sSFR shows a negative correlation with stellar mass. The star formation in more massive galaxies, although with higher SFR, has less impact than in less massive ones, due to the large stellar mass already formed. The same trend is observed in the local Universe, though shifted to lower sSFRs. This is in good agreement with the {\\em downsizing} scenario, in which massive galaxies are formed earlier than less massive ones. The fraction of massive galaxies (M$_{*}>$10$^{11}$M$_{\\odot}$) undergoing strong star formation processes ({\\em b}$>$2), $\\sim$29\\% at z$\\sim$0.84 against $<$3\\% at the local Universe, also supports this scenario.  Finally, we have quantified the {\\em downsizing} estimating the {\\em quenching} mass at z$\\sim$0.84 and at the local Universe based on the H$\\alpha$ star formation rate. We find that M$_{Q}$=1.0$^{+0.6}_{-0.4}\\times$10$^{12}\\, M_{\\odot}$ (log~(M$_{Q}$/M$_{\\odot}$)=12.0$\\pm$0.2) at z$\\sim$0.84 and M$_{Q}$=7.9$^{+1.9}_{-1.5}\\times$10$^{10}\\, M_{\\odot}$ (log (M$_{Q}$/M$_{\\odot}$)=10.9$\\pm$0.1) in the local Universe. The evolution since the local Universe is out of doubt, with an increase in the {\\em quenching} mass of an order of magnitude."
    },
    "1107/1107.3551_arXiv.txt": {
        "abstract": "Despite decades of study, it remains unclear whether there are distinct radio-loud and radio-quiet populations of quasi-stellar objects (QSOs).  Early studies were limited by inhomogeneous QSO samples, inadequate sensitivity to probe the radio-quiet population, and degeneracy between redshift and luminosity for flux-density-limited samples.  Our new 6 GHz EVLA observations allow us for the first time to obtain nearly complete (97\\%) radio detections in a volume-limited color-selected sample of 179 QSOs more luminous than $M_i = -23$ from the Sloan Digital Sky Survey (SDSS) Data Release Seven in the narrow redshift range $0.2 < z < 0.3$.  The dramatic improvement in radio continuum sensitivity made possible with the new EVLA allows us, in 35 minutes of integration, to detect sources as faint as 20 $\\mu$Jy, or $\\log[L_\\mathrm{6\\,GHz}({\\rm W\\,Hz^{-1}})]\\approx21.5$ at $z = 0.25$, well below the radio luminosity, $\\log[L_6({\\rm W\\,Hz}^{-1})]\\approx22.5$, that separates star-forming galaxies from radio-loud active galactic nuclei (AGNs) driven by accretion onto a super-massive black hole.  We calculate the radio luminosity function (RLF) for these QSOs using three constraints: (a) EVLA 6 GHz observations for $\\log[L_6({\\rm W\\,Hz^{-1}})]<23.5$, (b) NRAO-VLA Sky Survey (NVSS) observations for $\\log[L_6({\\rm W\\,Hz^{-1}})]>23.5$, and (c) the total number of SDSS QSOs in our volume-limited sample.  We show that the RLF can be explained as a superposition of two populations, dominated by AGNs at the bright end and star formation in the QSO host galaxies at the faint end. ",
        "introduction": "\\label{sec:intro}  Since the discovery of optically selected QSOs \\citep{sandage65}, the difference between the ``radio-loud\" (RLQ) and ``radio-quiet\" (RQQ) QSO populations remains elusive.  There has been a continuing controversy as to whether the radio luminosity distribution of QSOs is bimodal \\citep[e.g.,][]{kellermannEtal89,i02,i04qso} or is merely broad and smooth \\citep[e.g.,][]{cirasuoloEtal03,lacyEtal10}.  A bimodal distribution suggests that two distinct physical processes are present, with one process being significantly more powerful than the other.  The two primary sources of radio emission from galaxies are (1) accretion onto super-massive black holes in active galactic nuclei (AGNs) and (2) star formation.  Radio-AGN emission is due to synchrotron from relativistic plasma jets, and the associated hotspots and lobes resulting from jet interaction with the surrounding medium.  Star formation yields free-free emission from HII regions and synchrotron radiation from relativistic electrons believed to be accelerated in supernova remnants.  Radio AGNs can be extremely luminous, with radio emission reaching $\\log[L_\\mathrm{6\\,GHz}({\\rm W\\,Hz}^{-1})]\\approx27$ in the case of 3C~273 \\citep{3Cspectra}, whereas star formation leads to radio luminosities of $\\log[L_\\mathrm{6\\,GHz}({\\rm W\\,Hz}^{-1})]\\approx21$ for a galaxy like the Milky Way, or as high as $\\log[L_\\mathrm{6\\,GHz}({\\rm W\\,Hz}^{-1})]\\approx23.5$ in the case of a star-bursting galaxy such as Arp 220.  Radio-AGNs dominate the radio luminosity function (RLF) of galaxies brighter than $\\log[L_\\mathrm{6\\,GHz}({\\rm W\\,Hz}^{-1})]\\approx22.5$ in the radio, while star-formation in galaxies without an AGN dominates at fainter luminosities \\citep{condonEtal02}.  The question as to the nature of radio emission in radio-quiet QSOs has not been previously addressed using a homogeneous and complete QSO sample that is sensitive to luminosities significantly fainter than $\\log[L_\\mathrm{radio}({\\rm W\\,Hz}^{-1})]=22.5$.  It is possible that some, if not all, QSOs are hosted in star-forming galaxies, which only become visible at radio wavelengths when the AGN radio emission is below some given threshold. In this Letter, we use results from the Expanded Very Large Array (EVLA) \\citep{introPaper} to constrain the radio luminosity function (RLF) of QSOs. We show that the RLF is consistent with the hypothesis that QSO radio sources with 6 GHz spectral luminosity $\\log[L_6({\\rm W\\,Hz}^{-1})]>{23}$ are powered primarily by AGNs, while those with $\\log[L_6({\\rm W\\,Hz}^{-1})]<23$ are powered primarily by star formation in their host galaxies.  By taking advantage of recent advances in both radio and optical capabilities, we have, for the first time, obtained nearly complete radio detections in a large volume-limited sample of optically  selected QSOs.  In this Letter, we describe early results of our EVLA observations.  In Section~\\ref{sec:data} we describe the target selection and radio observations.  In Section~\\ref{sec:results} we present our results and interpret them in terms of the superimposed radio contributions from AGNs and star-forming host galaxies. Our conclusions are summarized in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  For the first time, our observations adequately sample the radio-quiet population; that is, they detect nearly all optically selected QSOs in a volume-limited sample by reaching $\\log[L_6({\\rm W\\,Hz^{-1}})]=21.5$.  Earlier studies did not have the sensitivity needed to study the full range of radio-quiet QSOs by reaching faint radio flux densities ($S\\ll1$\\,mJy) for optical samples with significant numbers of low-redshift ($z<0.5$) QSOs to probe the low end of the RLF.  Analyses based on radio-selected samples are inherently biased toward the radio-loud population, but previous studies reporting smooth luminosity distributions did not extend sufficiently faint to study the radio-quiet population.  The 6 GHz RLF of low-redshift color-selected QSOs is constrained by EVLA detections of sources stronger than $20\\,\\mu$Jy and by the very small fraction (6/179) of nondetections in the SDSS sample.  The strong sources constrain the AGN contribution above $\\log[L_6({\\rm W\\,Hz}^{-1})]\\approx23.5$.  The fainter sources imply a second radio contributor narrowly peaked around $\\log[L_6({\\rm W\\,Hz}^{-1})] \\approx 22.4$ and confirm the two-population model of QSO flux densities found by \\citet{kellermannEtal89}. Using the RLFs of nearby galaxies as examples, we suggest that this second contributor is radio emission produced by strong star formation in the QSO host galaxies, most of which therefore cannot be ``red and dead'' massive ellipticals.  In this Letter, we have presented the preliminary results of our investigation to characterize the radio population of QSOs using an optically selected, volume limited sample. We suggest that the RLF is a superposition of radio emission from AGNs and from host star formation. It will be important to better understand the relation between QSO emission in different wavelength regimes by characterizing the host galaxies of these QSOs using optical and infrared colors and morphology, as well as the radio morphology. The full details of this study will be presented in a separate paper."
    },
    "1107/1107.1835_arXiv.txt": {
        "abstract": "The Dispersed Fixed Delay Interferometer (DFDI) method is attractive for its low cost, compact size and multi-object capability in precision radial velocity (RV) measurements. The phase shift of fringes of stellar absorption lines is measured and then converted to an RV shift via an important parameter, phase-to-velocity scale (PV scale) determined by the group delay (GD) of a fixed delay interferometer. Two methods of GD measurement using a DFDI Doppler instrument are presented in this paper: 1), GD measurement using white light combs generated by the fixed delay interferometer; 2), GD calibration using an RV reference star. These two methods provides adequate precision of GD measurement and calibration given current RV precision achieved by a DFDI Doppler instrument. They can potentially be used to measure GD of an fixed delay interferometer for sub-meter precision Doppler measurement with a DFDI instrument. Advantages and limitations of each method are discussed in details. The two methods can serve as standard procedures of PV scale calibration for DFDI instruments and cross checks for each other. ",
        "introduction": "As of Apr 2012, there are over 700 discovered exoplanets, and most of them are detected by the radial velocity (RV) technique\\footnotemark\\footnotetext{http://exoplanet.eu/; http://exoplanets.org/}. RV precision of 1 $\\rm{m}\\cdot\\rm{s}^{-1}$ has been routinely achieved~\\citep{Bouchy2009,Howard2010} with instruments such as HARPS~\\citep{Mayor2003} and HIRES~\\citep{Vogt1994}, which are cross-dispersed echelle spectrographs. While cross-dispersed echelle spectrographs are commonly used in instruments for precision RV measurements, a method using a dispersed fixed delay interferometer (DFDI) has offered an alternative method~\\citep{Ge2006,Fleming2010,Lee2011}. In this method, a Michelson-type interferometer is used in combination with a moderate resolution spectrograph, RV signals are then extracted from phase shift of interference fringes of stellar absorption lines~\\citep{Erskine2000,Ge2002b,Ge2002,Erskine2003}. The details about the DFDI theory and applications are discussed in~\\citet{vanEyken2010} and~\\citet{Wang2011}. Instrument adopting the DFDI method has demonstrated advantages such as low cost, compact size and multi-object capability~\\citep{Ge2002,Ge2006,Fleming2010,Lee2011,Wisniewski2012}.  In the DFDI method, a fixed delay interferometer ~\\citep{Wan2009b,Wan2011} plays a crucial role in creating stellar spectral fringes  for high precision RV measurements~\\citep{Ge2002,Erskine2003}. The Doppler sensitivity can be optimized by carefully choosing the group delay (GD) of the interferometer~\\citep{Wang2011}. More specifically, GD of an interferometer should be chosen such that the spatial frequency of white light combs (WLCs) matches with that of a stellar spectrum after rotational broadening. GD is defined by the following equation: \\begin{equation} \\label{eq:GD_def} \\rm{GD}(\\nu)=-\\frac{1}{2\\pi}\\cdot\\frac{d\\phi}{d\\nu}, \\end{equation} where $\\phi$ is phase shift and $\\nu$ is optical frequency. The interferometer in a DFDI instrument is usually designed to be field-compensated to minimize the influence of input beam instability~\\citep{Wan2009b,Wang2010}. It is realized by carefully selecting glass materials and thicknesses of two second surface mirrors such that their virtual images are overlapped. Because glasses are used in the optical paths,  $\\phi$ does no longer linearly change with frequency, therefore GD is dependent on optical frequency. An inaccurate GD measurement may significantly limit the RV measurement accuracy~\\citep{Barker1974,vanEyken2010}.  In practice, there may be several methods of measuring GD: \\begin{enumerate} \\item Calculate GD based on glass refractive indices using Sellmeier equation and thicknesses from manufacturer specification. \\item Forward model the spectrum of a known spectral source, such as an Iodine cell or a Th-Ar lamp. \\item Measure phase and frequency using a while light source, such as a tungsten lamp. \\item Calibrate GD using a source with known velocity. \\end{enumerate}  Method 1 is straightforward but may lack of adequate precision because of uncertainty in parameters in Sellmeier equation and manufacturer tolerance for glass thickness. Method 2 holds great promise for accurately determining GD but there some current practical issues preventing us from adopting this method (see more detailed discussion in \\S \\ref{sec:WLCdisc}). We will use Method 3 and 4 to measure GD of an interferometer in this paper.  In the DFDI method,  GD determines the phase-to-velocity (PV) scale, the proportionality between the measured phase shift and the velocity shift. Since the DFDI method is realized by coupling a fixed delay interferometer with a post-disperser, the resulting fringing spectrum$-$stellar absorption lines superimposing on the WLCs$-$is recorded on a CCD detector (illustrated in Fig. \\ref{fig:DFDI_setup}). The fringe phase is expressed by the following equation: \\begin{equation} \\label{eq:fring_phase} \\phi(\\nu,y)=\\frac{2\\pi\\cdot\\tau(\\nu,y)\\cdot\\nu}{c}, \\end{equation} where $y$ is the coordinate along slit direction, which is transverse to dispersion direction, $\\tau$ is the optical path difference (OPD) of an interferometer and $c$ is the speed of light. Two mirrors (arms) of the interferometer are designed to be tilted towards each other along the slit direction such that several fringes are formed along each $\\nu$ channel. The intersection of a stellar absorption line and a WLC moves (from $P_o$ to $P$ in Fig. \\ref{fig:DFDI_setup}) if there is a shift of an absorption line due to a change of stellar RV. Consequently, a small change of $\\phi$ in the dispersion direction, $\\Delta\\phi_x$, is induced: \\begin{eqnarray} \\label{eq:fring_phase_shift} \\Delta\\phi_x & = & \\frac{d\\phi}{d v}\\cdot\\Delta v=\\frac{d\\phi}{d\\nu}\\cdot\\frac{d\\nu}{d v}\\cdot\\Delta v \\nonumber \\\\ & = & \\frac{d\\phi}{d\\nu}\\cdot\\frac{\\nu}{c}\\cdot\\Delta v=\\Gamma\\cdot\\Delta v, \\end{eqnarray} where $\\Gamma$ is defined as phase-to-velocity scale (PV scale). It is determined by the GD of an interferometer, which becomes explicit if Equation (\\ref{eq:GD_def}) and (\\ref{eq:fring_phase_shift}) are combined: \\begin{equation} \\label{eq:PVS} \\Gamma=-2\\pi\\cdot\\rm{GD}\\cdot\\frac{\\nu}{c}. \\end{equation}  At resolutions typically adopted by the DFDI method ($5,000\\le R\\le20,000$), stellar lines (line width$\\sim$0.1$\\AA$) are not resolved and a measurement of $\\Delta\\phi_x$ is extremely difficult. Instead, $\\Delta\\phi_y$, phase shift along $y$ direction can be measured, which is equal to $\\Delta\\phi_x$ if the combs generated by an interferometer are parallel to each other. This is a good approximation at very high orders of interference. The advantage of measuring $\\Delta\\phi_y$ instead of $\\Delta\\phi_x$ is seen from Fig. \\ref{fig:DFDI_setup}, in which the physical shift in the $\\nu$ direction is amplified in $y$ direction, the amplification rate is determined by the relative angle between the interferometer combs and a stellar absorption line. Therefore, $\\Delta\\phi_y$ is relatively easier to measure compared to $\\phi_x$ and it is measured by fitting a well-sampled periodical flux signal along the $y$ direction in the DFDI method. Compared to conventional high-resolution Echelle method, the number of freedom for the DFDI method in the fitting process is much less and small Doppler phase shift can be relatively easier detected with a simple functional form, i.e., a sinusoidal function. However, we want to point out that while the DFDI method provides a boost in instrument Doppler sensitivity, the Doppler sensitivity is not strongly dependent on the amplification rate because flux slope decreases as amplification rate increases, which negates the gain of phase slope.  The paper is organized as follows: in \\S \\ref{sec:OpdMeasurement}, we present a new method of GD measurement using a DFDI Doppler instrument, which exploits WLCs generated by a fixed delay interferometer. In \\S \\ref{sec:OPD_solution}, we present another method of GD calibration using an RV reference star. Observation results after implementing the newly measured GD are presented in \\S \\ref{sec:Implementation}. In \\S \\ref{sec:Conclusion}, we summarize and discuss the new results in this paper. ",
        "conclusions": "\\label{sec:Conclusion}  \\subsection{Summaries}  The PV scale is an important parameter in the DFDI method that translates a measured phase shift to an RV shift, and is determined by the group delay (GD) of an interferometer. We have provided and discussed two methods of GD measurement and calibration: 1), GD measurement using white light combs (WLCs) generated by the interferometer in a DFDI Doppler instrument; 2), GD calibration using an RV reference star (RS). Table \\ref{tab:Cal_Comp} summarizes the main results and the comparison between these two methods. The accuracy of GD measurement is sufficient for current RV precision achieved with instruments using the DFDI method~\\citep{Fleming2010, Lee2011, Muirhead2011}. However, higher measurement and calibration precision is required in the near future as higher RV precision is achieved by DFDI instruments in search for exoplanets. RS and WLC methods can serve as complementary methods of GD measurement and calibration for DFDI instruments.  \\subsection{Discussions} \\subsubsection{White Light Comb (WLC) Method} \\label{sec:WLCdisc}   The GD measurement using WLCs created by the interferometer provides a direct way of calibrating the PV scale. In the region where combs are visible, effective S/N is relatively low ($\\sim$15) because of  low comb visibility (1.5\\%). In addition, GD cannot be measured in the region where combs are not visible, which limits the application of this method. We are able to measure GD in a region that accounts for half of the spectrum coverage. Extrapolation beyond the measurement range may result in large uncertainties. The major issue facing the method is that the data reduction pipeline may have introduced unknown errors while correcting optical distortions such as spectrum curvature and slant.  In principle, we can use a tungsten lamp with an iodine cell or a Th-Ar lamp instead of a tungsten lamp in order to increase the fringe visibility. However, there are some practical concerns that hinder us from applying the above solutions: 1), line blending, because of low spectral resolution, many spectral lines can not be resolved and it is not certain at this stage how line blending affects phase measurement; 2), illumination correction, which is required to correct for illumination profile in slit direction in order to properly measure the phase. We adopt a self-illumination correction procedure in the pipeline which requires a certain continuum level to be successfully achieved. Th-Ar is less affected by line blending if a careful line selection process is involved, but it does not have enough continuum level for self-illumination correction. An experiment is being conducted in which a second interferometer is used to improve the visibility of WLCs so that GD is more precisely measured at a higher effective S/N for a wider frequency coverage.  When compared to previous work in the field of GD measurement,~\\citet{Amotchkina2009} achieved a measurement precision of $1\\times10^{-4}$ ps. Our measurement of GD has a typical accuracy of $\\sim6\\times10^{-3}$ ps (limited by effective S/N and systematic errors), which is more than an order of magnitude lower. However, it is shown in \\S \\ref{sec:Err_Ana} that substantial improvement would be able to be achieved once the data reduction pipeline has a better handle of optical distortion.~\\citet{Wan2010} measured GD for the MARVELS interferometer using a scanning WLI method and achieved a precision of $0.6\\times10^{-5}\\sim2.4\\times10^{-5}$ ps, which is more than two order of magnitude better than the results in this paper. However, there are practical concerns using GD measurement results from the scanning WLI method because they are not measured in situ, therefore it is not easy to associate a position in a WLI measurement to a fiber position.    \\subsubsection{Reference Star (RS) Method}  The GD calibration using an RV reference star (RS) is a self-calibrating process and has the potential of achieving a high calibration precision if the following requirements are met: 1), the RV reference star has a large velocity shift during a observation window; 2), the RV reference star is bright; 3), the data reduction pipeline is able to produce the photon-limited RV precision. In addition, GD is practically measured within a certain band width: \\begin{equation} \\label{eq:GD_cal_ref_wb} \\rm{GD}(\\nu)=\\frac{\\int_{\\Delta\\nu}\\rm{GD}(\\nu)\\omega(\\nu)d\\nu}{\\int_{\\Delta\\nu}\\omega(\\nu)d\\nu}, \\end{equation} where $\\omega(\\nu)$ is weight function. The band width, $\\Delta\\nu$, should be small such that the dispersion effect is negligible. The limitations stated above prevent us from precisely determining the PV scale at the position of each fiber because not every fiber has a continuous observation on a bright known RV reference star. For MARVELS, the brightest RV reference star available has V mag of 8 and the resulting S/N is $\\sim$100 per pixel. However, the RS method is a very promising approach for a single-object DFDI instrument because only one bright reference star is required in the field. We are planning to apply this method in calibrating GD of another DFDI instrument (EXPERT) at KPNO 2.1m telescope~\\citep{Ge2010}. Note that the S/N can be further increased by conducting multiple independent measurements and increasing instrument throughput.   We would like to express our deepest gratitude to the anonymous referee of this paper, without whom there would not have been this paper. We  acknowledge the support from NSF with grant NSF AST-0705139, NASA with grant NNX07AP14G (Origins), UCF-UF SRI program, DoD ARO Cooperative Agreement W911NF-09-2-0017, SDSS III consortium, Dharma Endowment Foundation and the University of Florida.  Funding for SDSS-III (\\url{http://www.sdss3.org/}) has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science.  SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, University of Cambridge, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck Institute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University."
    },
    "1107/1107.2946.txt": {
        "abstract": "We examine atmospheric heating by radio active galactic nuclei (AGN) in distant X-ray clusters by cross correlating clusters selected from the 400 Square Degree (400SD) X-ray Cluster survey with radio sources in the NRAO VLA Sky Survey.  Roughly 30\\% of the clusters show radio emission above a flux threshold of 3 mJy within a projected radius of 250 kpc.  The radio emission is presumably associated with the brightest cluster galaxy.  The mechanical jet power for each radio source was determined using scaling relations between radio power and cavity (mechanical) power determined for nearby clusters, groups, and galaxies with hot atmospheres containing X-ray cavities.  The average jet power of the central radio AGN is approximately $2\\times 10^{44}$\\ergs. We find no significant correlation between radio power, hence mechanical jet power, and the X-ray luminosities of clusters in the redshift range 0.1 -- 0.6.  This implies that the mechanical heating rate per particle is higher in lower mass, lower X-ray luminosity clusters.  The jet power averaged over the sample corresponds to an atmospheric heating of approximately 0.2 keV per particle within R$_{500}$.  Assuming the current AGN heating rate does not evolve but remains constant to redshifts of 2, the heating rate per particle would rise by a factor of two.  We find that the energy injected from radio AGN contribute substantially to the excess entropy in hot atmospheres needed to break self-similarity in cluster scaling relations.  The detection frequency of radio AGN is inconsistent with the presence of strong cooling flows in 400SD clusters, but does not exclude weak cooling flows.  It is unclear whether central AGN in 400SD clusters are maintained by feedback at the base of a cooling flow.  Atmospheric heating by radio AGN may retard the development of strong cooling flows at early epochs. ",
        "introduction": "\\label{sec:intro}   Heating by active galactic nuclei (AGN) may be responsible for two significant phenomena in  galaxy clusters. First, cool, dense cores found in the X-ray atmospheres of clusters  are expected to cool efficiently to low temperatures \\citep[][]{fabian94}.  However,  the levels of cold gas and star formation lie well below the values expected from pure radiative cooling \\citep[reviewed by][]{peterson06, mcnamara07}.  The generally accepted solution is that the cooling is being compensated by one or more heating mechanisms, possibly including conduction from the hot gas in non-cooling regions \\citep[e.g.][]{narayan01, zakamska03, voigt04, wise04}, and AGN heating \\citep[e.g.][]{binney95, churazov02, roychowdhury04, voit05}. Among the possible heating mechanisms, AGN heating stands out because of both the enormous energy released and the opportunity for feedback, since AGN power is controlled by the accretion rate.  Second, galaxy groups and poor clusters are less luminous than expected based on the $L_{\\rm X}$-$T$ relation derived for massive clusters \\citep{markevitch98, arnaud99, ruszkowski04, sanderson03,popesso05}. This is interpreted as a break of the self-similar scaling relation between X-ray luminosity and gas temperature expected in structure formation models that include only gravitational heating \\citep[][]{kaiser86,evrard96}.  The lower than expected X-ray luminosity for a given temperature is thought to reflect extra entropy (heat) in the gas that was injected by star formation and AGN at early times \\citep[e.g.][]{kaiser91, balogh99, borgani05, ponman03, croston05, sanderson05, jetha07, cavagnolo09}.   The ``preheating\"  model proposes that the excess entropy was injected during the early stages of structure formation \\citep[e.g.][]{kaiser91, balogh99, mccarthy02}.  The amount of energy released by an AGN, on the order of $10^{62}\\,M_{\\rm BH} / [{10^{9}\\,{\\rm M}_{\\odot}}]$\\,erg, is more than enough to be the major heating source. Nevertheless, how and where the energy is distributed into the ICM remains an issue. As discussed by \\citet{mcnamara07}, AGN heat in two ways.  First, AGN outbursts inflate cavities.  The enthalpy released as the cavities rise buoyantly is dispersed into the ICM \\citep{bruggen02, reynolds02, churazov02, birzan04}.  Thermalization of the energy of cosmic rays diffusing out of a radio lobe is an alternative channel for converting lobe free energy to heat \\citep[e.g.][]{mathews09}. Second, the AGN inject energy into the ICM through shocks \\citep[e.g.][]{forman05, mcnamara05, nulsen07} and sound waves \\citep[e.g.][]{fabian05}.  Cavity power can be estimated using the enthalpy contained in the cavity and its rise time. In practice, the enthalpy is calculated as ${4pV_{\\rm cav}}$, assuming the cavity is filled with relativistic particles, and the age can be estimated using the terminal velocity of the buoyantly rising bubbles \\citep[][]{birzan04, dunn05, rafferty06, birzan08}. All of these quantities, the pressure of the surrounding gas $p$, the volume of the cavity $V_{\\rm cav}$ and sound speed in the ICM, can be obtained from X-ray data in principle. Simulations by \\citep[][]{mendygral11} have verified that the cavity power estimated using this approach can be accurate to approximately a factor of two. Furthermore, \\citet{birzan08} and \\citet{cavagnolo10} find that the radio power and the buoyant cavity power are correlated for samples ranging from elliptical galaxies and galaxy groups to massive galaxy clusters. Thus, the radio power can be used as a proxy for the cavity power for large samples of radio galaxies embedded in hot atmospheres that lack deep X-ray observations. However, cavity power estimated using this approach is probably a lower limit, because it does not include the additional energy released by shocks \\citep[][]{forman05,mcnamara05,nulsen05a,nulsen05b}.  It is more challenging to calculate the total AGN energy deposited over the life of a cluster.  Since the power output of AGN is highly variable and outbursts are intermittent, current AGN power is not necessarily representative of its historical average.  Even under the simplest assumption, that the power of an AGN is the same in every outburst \\citep[which is inconsistent with data, e.g.][]{fabian06,randall11}, an estimate is still required for the fraction of the time that the AGN is in outburst, i.e., for its duty cycle, which is also inaccessible for individual AGN.  An alternative to studying single systems is to determine the average AGN power for a large sample.  For a well chosen sample, the average power should represent the time average of the AGN power for individual sample members.  The key to success here is to use a well-defined, unbiased, and sufficiently large sample.  In this paper, we estimate the average jet power in distant, X-ray selected clusters found in the 400 Square Degree  Cluster Survey \\citep[400SD;][]{burenin07,vikhlinin05} by cross correlating the 400SD catalog with the  NRAO VLA Sky Survey \\citep[NVSS;][]{nvss}.  We calculate jet powers for radio sources above a flux limit of $\\sim3$\\,mJy.  Following earlier work on AGN heating in cooling core clusters \\citep[][]{dunn06,dunn08,rafferty06,rafferty08}, we evaluate the AGN power injected into normal clusters to address the question of how much energy is deposited by AGN into cluster halos over a significant fraction of their ages.  Note that AGN feedback would be expected to enhance the incidence of radio outbursts at the centers of cooling core clusters.  As a result, this question cannot be answered based on a sample of nearby cooling core clusters, which does not represent the general cluster population. Also, since we are interested in the total power injected by AGN, multiple radio sources may be attributed to a single cluster in the present study.  The paper is arranged as follows. In \\S\\ref{sec:data}, we introduce the 400SD clusters, and the matched radio sources. In particular, we discuss the background source correction in \\S\\ref{sec:data_nvssbackground}, and examine the morphology of a subsample of the NVSS sources using images from the survey Faint Images of the Radio Sky at Twenty-cm \\citep[FIRST; ][]{first} in \\S\\ref{sec:data_first}. Then we present our analysis and discuss the interpretation: the correlation between the 1.4\\,GHz power of the radio AGNs and the X-ray luminosity of their host clusters in \\S\\ref{sec:power_correlation}, the radial distribution of radio sources in the clusters in \\S\\ref{sec:distribution_radio}, the jet power estimated from the radio power in \\S\\ref{sec:cavitypower}, and the estimation of AGN heating energy in \\S\\ref{sec:agnheating}. We summarize our results in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:result}   \\subsection{Correlation between radio power of AGN and X-ray luminosity of clusters?}\\label{sec:power_correlation}  Using our catalog of 400SD clusters, we now examine the relationship between the radio power of the central galaxies and cluster X-ray luminosity.  A simple flux-flux plot in Fig.~\\ref{fig:correlation_flux} shows no correlation between the radio fluxes\\footnote{Radio fluxes are summed for all sources within 250\\,kpc of each cluster cener. However, most of the clusters (51 out of 56) with radio sources within 250\\,kpc radius only contain a single radio source, so that the fluxes could be taken as those of individual radio sources.} (F$_{1.4}$) of the sources within $250$\\,kpc radius and the X-ray fluxes (F$_{\\rm X,0.5-2.0keV}$). The Spearman correlation is only $0.15\\pm0.20$.  To make sure the correlation is not hidden by redshift dependent factors, such as reddening, we plot the ratio of radio power to X-ray luminosity versus redshift in Fig.~\\ref{fig:correlation_powerratio}.  Although the points seem to be centered on a constant ratio of $\\sim 0.001$, the lower limits, estimated for clusters with no detected NVSS sources (F$_{1.4}>3$\\,mJy), and the upper limit (red curve) for clusters not detected at the X-ray flux limit of the 400SD survey, suggest that any correlation is deceptive.  In summary, no correlation is found between radio power and X-ray luminosity in our cluster sample.  Several previous studies \\citep[e.g.][]{sun09,mittal09} have demonstrated that the correlation between X-ray luminosity of the cooling core and the radio power only appears in ``strong\" cooling core clusters.  Therefore, the lack of a correlation between radio and X-ray power is consistent with the lack of strong cooling core clusters among 400SD clusters \\citep{vikhlinin07,santos10,samuele11}. Of the 52 400SD clusters with Chandra data, 20 have NVSS sources within 250 kpc of the center.  Of those, 6 clusters have surface brightness concentrations $c_{\\rm SB} > 0.1$ \\citep[][]{santos08}, corresponding to central cooling times shorter than $\\simeq 5$ Gy \\citep[][]{santos10}.  These are likely to be weak cool core clusters \\citep[WCC in ][]{mittal09}. These sources are plotted in red in Fig.~\\ref{fig:correlation_flux} and their radio and X-ray fluxes do appear better correlated than for the bulk of the sample, consistent with previous studies.   We also consider the dependence of the fraction of clusters with NVSS radio sources on X-ray luminosity. In Fig.~\\ref{fig:dutycycle}, the incidence of radio sources appears to increase slightly with X-ray luminosity.  However, the range of X-ray luminosities sampled depends on the redshift in a flux-limited survey. The low luminosity clusters in the sample are only detectable at lower redshift because of the X-ray flux limit.  Separating the sample into two redshift slices ($0.1< z <0.3$ and $0.3< z <0.5$ in Fig.~\\ref{fig:dutycycle}), we find that the apparent correlation between X-ray luminosity and the fraction of cluster with NVSS sources disappears.  Thus, we suggest that the difference between the blue points (higher redshift, more luminous clusters) and the red points (lower redshift, less luminous clusters) is probably caused by redshift and not X-ray luminosity. \\citet{best07a} found no correlation between the fraction of radio-loud AGN in cluster galaxies and cluster velocity dispersion, which is, a proxy for X-ray luminosity and cluster mass.  However, \\citet{lin07} found a clear increase in the ``radio active fraction\" in a massive cluster sample selected when they have bright BCGs with ${\\rm M}_{k}<-24$.  In order to examine this apparent inconsistency, we need to study the k-band luminosity of the BCGs which is beyond the scope of this paper.  Although we see no evidence of a dependence on X-ray luminosity, the incidence of radio AGN does increase marginally at higher redshift. This may reflect higher AGN fractions in the past, as suggested by some, e.g.,by \\cite{martini07, martini09} for X-ray AGN and by \\citet{galametz09} for AGN selected by radio, infrared and X-ray. However, \\citet{gralla10} analyzed radio AGN in the 618 clusters of the Red Sequence Clusters Survey at redshifts $0.35< z <0.95$, finding that the average number of AGN is independent of redshift.  Although they are less significant, our results agree better with the trend found for radio AGN by \\citet{galametz09}, of an increase in incidence by about a factor of 2 ($\\sim 2\\sigma$ significance) from $z<0.5$ to $0.5<z<1.0$.  \\begin{figure}[hb] %fig7 %\\epsscale{1.3} \\vspace*{-5mm} \\hspace*{-7mm}\\plotone{correlation2_0250kpc.ps} \\figcaption[]{$F_{1.4}$ versus $F_{\\rm X}$ for the 400SD clusters with NVSS sources within $250$\\,kpc.  The expected background fluxes (\\S\\ref{sec:data_nvssbackground}) have been subtracted, although the correction is insignificant for this small region. Using the archived Chandra data, we can estimate the cooling time of clusters referring from the concentration parameter $c_{SB}$. The clusters with cooling time less than 5\\,Gyr are plotted in red. \\label{fig:correlation_flux}} \\end{figure}  \\begin{figure}[hb] %fig8 %\\epsscale{1.3} \\vspace*{-5mm} \\hspace*{-7mm}\\plotone{corrNLUM_0250kpc.ps} \\figcaption[]{Ratio of radio power within 250 kpc to X-ray luminosity for 400SD clusters.  Dots mark clusters with detected radio sources and upper limits are for non-detections, assuming one 3\\,mJy radio source in each cluster.  The red curve shows the X-ray flux limit of the survey, for the median radio flux. \\label{fig:correlation_powerratio}} \\end{figure}    \\begin{figure}[h] %fig9 %\\epsscale{1.3} \\vspace*{-5mm} \\hspace*{-7mm}\\plotone{400sd_numsrcs_lx.ps} \\figcaption[]{Fraction of radio AGN in cluster cores versus X-ray luminosity.  The ratio of the number of clusters with radio loud AGN within 250\\,kpc) to the total number of clusters in each X-ray luminosity bin is plotted against the X-ray luminosity.  To reveal the effect of the redshift-dependent luminosity cut, the clusters are separated into two subsamples by redshift, $0.3< z <0.5$ (blue) and $0.1< z <0.3$.  The mean redshift of each bin is marked on the plot.  Only radio sources with P$_{1.4}>3\\,10^{24}$\\,W\\,Hz$^{-1}$ are counted, corresponding to the radio power limit at ${\\rm z}\\sim0.5$.  Uncertainties are calculated assuming Poisson statistics. \\label{fig:dutycycle}} \\end{figure}   \\begin{figure}[h] %fig10 %\\epsscale{1.3} \\vspace*{-5mm} \\hspace*{-7mm}\\plotone{400sd_numsrcs_radius.ps} \\figcaption[]{Distribution of radio sources around 400SD clusters. The radio sources are grouped into two subsamples by redshift of the host cluster: $0.1< z < 0.3$ for the solid line and $0.3< z < 0.5$ for the red dashed line.  The gray solid line shows the distribution for lower redshift sources, for clusters with $L_{\\rm X,bol}>10^{44}$\\,erg\\,s$^{-1}$ for the host clusters.  The numbers have been corrected for background sources (\\S\\ref{sec:data_nvssbackground}).  Error bars show the standard deviations for each point, while those for the gray solid line are omitted for clarity.\\label{fig:400sd_numsrcs}} \\vspace*{1cm} \\end{figure}   \\subsection{Radial Distribution and Incidence of radio AGN in clusters}\\label{sec:distribution_radio}  In Fig.~\\ref{fig:400sd_numsrcs}, we plot the average cumulative number of radio AGN within radius $r$ versus $r$ for the 400SD clusters.  The number of radio sources increases from an average of $\\sim0.3$ in the cluster cores to $\\sim 1$ within $2$\\,Mpc.  However, the number of radio sources increases much more slowly than the area (i.e., $\\propto r^2$), consistent with the observation that the surface density of radio sources peaks at cluster centers \\citep{ledlow95a, best07a, lin07, croft07}.  Since most 400SD clusters have no more than one radio source within 250 kpc, the average number of radio sources shown as the leftmost point in Fig.~\\ref{fig:400sd_numsrcs} is equivalent to the radio source fractions in Fig.~\\ref{fig:dutycycle}.  The average number of radio sources we find in the cores of 400SD clusters is significantly lower than the fraction of radio-loud BCGs in nearby cooling core clusters, which lies between 75\\% to 100\\% \\citep{burns90,edwards07, dunn06b,mittal09}. This fraction is expected to correlate with the central cooling time. Our detection fraction is close to the fraction of radio sources found in nearby, non-cooling core clusters ($45\\%$) as defined by  \\citet{mittal09},  and is identical to the fraction found in optically selected clusters (30\\%) determined by \\citet{best07a}. While a more quantitative comparison between the numbers is difficult, a low radio source fraction for the 400SD clusters is consistent with our claim that most of them lack strong cooling cores.  At the same time, our detection fraction is significantly higher than the radio loud fraction of isolated ellipticals, which is $\\sim 15\\%$ \\citet{best07a}. So it is possible that our host galaxies lie at the centers of weak cooling flows.  This issue will need to be addressed using deep X-ray imaging.  In Fig.~\\ref{fig:400sd_numsrcs}, the clusters are divided into two redshift ranges, $0.1< z <0.3$ (solid line) and $0.3< z <0.5$ (red dashed line).  We find that the average number of radio sources increases with redshift, consistent with the discussion of Fig.~\\ref{fig:dutycycle} in \\S\\ref{sec:power_correlation}. Nevertheless, the comparison between the higher redshift and the lower redshift clusters   shows that the spatial distribution of radio sources around these clusters is not significantly different from that for all lower redshift clusters in the 400SD sample. In order to show that the redshift dependence of the average number of radio sources is not a consequence of the X-ray flux limit (i.e., luminosity limit) for the higher redshift clusters, we also plot the distribution of radio sources for lower redshift clusters with $L_{\\rm X, bol} > 10^{44}\\rm\\,erg\\,s^{-1}$ (gray solid line). This X-ray luminosity corresponds approximately to the flux limit for the higher redshift sample, so the close correspondence between the gray and black curves shows that their difference from the radial distribution for the higher redshift sample (red dashed line) is not due to the X-ray flux limit.  \\begin{figure}[h] %fig11 %\\epsscale{1.3} \\vspace*{-5mm} \\hspace*{-7mm}\\plotone{400sd_accumpower_radius.ps} \\figcaption[]{Distribution of radio power around 400SD clusters.  The symbols are the same as for Fig.~\\ref{fig:400sd_numsrcs}, but here the average power of NVSS radio sources within $r$ is plotted against $r$.   Note that the brightest source (3C288, with P$_{1.4}\\sim 700\\times10^{24}$\\,W\\,Hz\\,$^{-1}$, see Fig.~\\ref{fig:nvsslum}) at ${\\rm z}\\sim0.25$ is omitted (see \\S\\ref{sec:distribution_radio}).  If its radio power were included, the average powers for the lower redshift sources would be increased by $5\\times10^{24}$\\,W\\,Hz$^{-1}$.  The uncertainty in the average power is calculated assuming a Gaussian distribution. \\label{fig:400sd_accumpower}} \\end{figure}  Like Fig.~\\ref{fig:400sd_numsrcs}, Fig.~\\ref{fig:400sd_accumpower} shows the distribution of the average radio power around 400SD clusters.  The radio power also show a modest increase with redshift, but here the impact of the X-ray flux limit (see the gray solid line) is more marked. Note that the average power for the low redshift sample does not include the brightest radio source (3C288, ${\\rm z}=0.246$), which is roughly an order of magnitude brighter than the second brightest source within $250$\\,kpc (Fig.~\\ref{fig:nvsslum}).  If the contribution of this source was included, the average power of the low redshift sample would exceed that for the high redshift sample.  We omit this source because the jet power derived from its radio power exceeds the jet power determined from its cavities by more than an order of magnitude \\citep[see \\S\\ref{sec:cavitypower}, ][]{lal10}. This discrepancy may simply reflect the large intrinsic scatter in the relationship of Eqn.~\\ref{eqn:high}, but we are unable to pursue this issue here, since our sample lacks a statistically significant population of such extremely luminous, short-lived sources.    \\begin{figure}[h] %fig12 %\\epsscale{1.3} %\\vspace*{-5mm} \\hspace*{-5mm}\\plotone{nprob_mc_z_combine.2.ps} \\figcaption[]{Top: Average jet power (\\meanpcav) of radio galaxies in the 400SD clusters.  We estimate the jet power of radio sources from the 1.4GHz flux in the NVSS catalogue, using the calibration of \\citet{cavagnolo10}. The averages are calculated in 5 redshift bins covering $0.1< z < 0.6$, for the matched NVSS sources within $250$\\,kpc (black circles) and 500\\,kpc (red circles) of each cluster center.  For the lowest  redshift bin, results are also shown for subsamples with ${\\rm L_{X}}>10^{44}$\\,erg\\,s$^{-1}$, plotted as stars.  The numbers of clusters hosting NVSS sources and the total numbers of clusters in each redshift bin are noted at the bottom/top of the plot.  The error bars, estimated using a Monte-Carlo method, are dominated by the scatter (\\shigh) in the correlation between jet power and radio luminosity.  Bottom: Ratio of jet power to X-ray luminosity for 400SD clusters versus redshift.  The plot shows the average jet powers from the top panel of Fig.~\\ref{fig:cavitypower} divided by the average total X-ray luminosity of the clusters in each redshift bin.  \\label{fig:cavitypower}} \\end{figure}  \\begin{figure}[h] %fig13 %\\epsscale{1.3} %\\vspace*{-5mm} \\hspace*{-5mm}\\plotone{nprob_mc_lum_cpDIVxlum4_addnosrc_400SD.2.ps} \\figcaption[]{Ratio of jet power to X-ray luminosity versus X-ray luminosity.  Similar to Fig.\\ref{fig:cavitypower}, but here the sample is grouped by X-ray luminosity. \\label{fig:cavitypower_lum_lum}} \\end{figure}     \\subsection{Jet power estimated from radio power}\\label{sec:cavitypower}  Radio power has been shown to correlate with ``cavity power,'' an estimate of the mean power of radio jets based on enthalpies and ages of X-ray cavities \\citep{birzan04, birzan08, cavagnolo10}.  Given their close correspondence to radio lobes, cavities are assumed to be formed by AGN outbursts.  As a cavity rises buoyantly and expands, its enthalpy is expected to be released into the hot atmosphere that hosts it \\citep[reviewed in][]{mcnamara07}.  This is one of several mechanisms by which radio outbursts may heat the gas.  Other mechanisms include shock waves or sound waves created by the expansion of cavities \\citep{forman05, mcnamara05,fabian05b, nulsen05a, nulsen05b, forman07} and thermalization of cosmic rays leaked from a radio source \\citep{bohringer88, mathews09,guo11}.  Here, we estimate jet powers, \\pcav, using the \\phigh\\ -- \\pcav\\ scaling relation of \\citet{cavagnolo10}, \\begin{eqnarray} \\log~\\pcav\\ = 0.75~(\\pm 0.14)~\\log~{\\phigh\\ } + 1.91~(\\pm 0.18),  \\label{eqn:high} \\end{eqnarray} which is expected to provide, on average, a minimum estimate for the total power injected into a cluster by a radio AGN.  Here, \\pcav\\ is in units of $10^{42}\\,{\\rm erg\\,s^{-1}}$, and \\phigh\\ in unit of $10^{40}\\,{\\rm erg\\,s^{-1}}$. The scatter in the correlation between jet power and radio luminosity is $\\shigh =0.78$\\,dex.  It is dominated by the intrinsic scatter \\citep[][]{birzan08}, due to a number of effects, including aging, and differences in particle content and magnetic field.  The relationship (Eqn.~\\ref{eqn:high}) is determined over 6 decades in \\phigh\\, ($10^{37}$-$10^{43}$\\,\\ergs), including data for nuclear radio sources for BCGs in cooling core clusters, with $L_{\\rm X}$ up to $10^{45}$\\,\\ergs, and low power sources in giant ellipticals with $L_{\\rm X}$ around $10^{43}$\\,\\ergs. The large scatter in this scaling relation implies that the jet power estimate it provides for any individual object has an uncertainty approaching a factor of 10.  We estimate the jet power for all of the cluster radio sources using Eqn.~\\ref{eqn:high}, albeit with a large intrinsic uncertainty. Assuming the errors are random, we expect the uncertainty in the ensemble average to be reduced to an acceptable level.  More importantly, the average jet power calculated for an unbiased sample should provide an estimate of the contribution to the energy budget of AGN power integrated across the sample.  A more precise measurement of jet power may be obtained for an individual object by analyzing its X-ray cavities.  Despite being more precise, the measurement is only an instantaneous snap shot of its current jet power, and does not necessarily reflect the historical average AGN power. The ages of FR 1 radio AGN outbursts in clusters lie typically between $10^{7}$ to $10^{8}$\\,yrs, and the lives of powerful FRII sources are even shorter \\citep{odea09}. These time scales are much less than cluster ages, which greatly exceed $10^9$ years. Therefore, an estimate of the time average jet power is required to calculate the total AGN energy input to the clusters over cosmic time.  This quantity may be estimated by averaging jet power over a sufficiently large and unbiased sample of clusters, as we have done here.  In Fig.~\\ref{fig:cavitypower}, we show our estimate of the average jet power (\\meanpcav), using a Monte-Carlo method that allows for the uncertainty in the radio flux and the parameters in Eqn.~\\ref{eqn:high}, the distribution of radio spectral indices\\footnote{A Gaussian distribution with a mean of 0.75 and a $\\sigma$ of 0.10 \\citep[][]{condon92} is assumed for the spectral indices.}, and the large intrinsic scatter (\\shigh) in the relation (Eqn.~\\ref{eqn:high}) to estimate confidence ranges.  Overall, the confidence ranges shown in Fig.~\\ref{fig:cavitypower} are dominated by \\shigh.  Note that we use the arithmetic mean value of $\\pcav$, excluding the brightest source (see \\S\\ref{sec:distribution_radio}).  While this gives a mean power that is substantially greater than the ``mean\" of the log-normal distribution of jet powers, assumed implicitly in the correlation (Eqn.~\\ref{eqn:high}), our purpose is to determine the average energy deposited by AGN, not its distribution, so that the arithmetic mean is the appropriate estimator.  In order to quantify the potential underestimate, we calculated the mean jet power excluding the 10\\% most powerful radio sources, and found that the number drops by 5\\%.    We find that the average jet power in the cores of the 400SD clusters is $2\\times10^{44}$\\ergs and does not show a significant redshift evolution. The slight increase in the jet power at higher redshift reflects the increased AGN fraction shown in Fig.~\\ref{fig:dutycycle} and Fig.~\\ref{fig:400sd_numsrcs}.  The typical total X-ray luminosity of the 400SD clusters (Fig.~\\ref{fig:xlum}, $0.2<{\\rm z}< 0.6$) is $\\sim3\\times10^{44}$\\,\\ergs, so that the average jet power in these clusters is about $50\\%$ of their X-ray luminosity (Fig.~\\ref{fig:cavitypower}).  Comparing the two plots in Fig.~\\ref{fig:cavitypower}, we find that the higher ratio of jet power to X-ray luminosity for the lowest redshift bin shown in the bottom panel of Fig.~\\ref{fig:cavitypower} is caused primarily by selection effects.  Using an X-ray flux limit reduces the average X-ray luminosity at lower redshifts, while lower luminosity clusters at higher redshifts fall below the detection threshold. Excluding low luminosity clusters ($L_{\\rm X}<10^{44}$\\ergs) for the low redshift bin $0.1< z < 0.2$ gives an average jet power (star symbols in Fig.~\\ref{fig:cavitypower}) that is comparable to the values at higher redshifts, consistent with the average radio power being largely independent of the redshift, as argued in \\S\\ref{sec:distribution_radio} (see Fig.~\\ref{fig:400sd_accumpower}). It suggests that the average jet power does not depend on the X-ray luminosity of a cluster.  If so, the impact of AGN on the gas will be greater in lower luminosity, hence less massive, systems. Fig.~\\ref{fig:cavitypower_lum_lum} shows that the minimum jet power may exceed the power radiated by the ICM in group-scale objects with $\\rm L_{x}<10^{44}$\\,\\ergs.  If the jet power is mostly retained by the hot atmospheres and not radiated away, AGN can cause a net increase of entropy in low mass systems.  %We note that the jet power of the extreme radio source 3C288 was %omitted from the calculations here (\\S\\ref{sec:distribution_radio}). %Its jet power of $9.8\\times10^{45}$\\,\\ergs estimated from the radio %power is $\\sim$15 times larger than the value estimated by %\\citet{lal10} ($6.1\\times10^{44}$\\,\\ergs).  The radio power may be %overestimated due to contributions of bright spots in the complex %radio structure shown in the top left panel of %Fig.~\\ref{fig:examples}.  In principle, the jet power should be %determined from the diffuse power emitted by the radio lobes alone, %but that cannot be separated at the resolution of the NVSS.  Although %this difference is consistent with the large scatter in the relation %(Eqn.~\\ref{eqn:high}), 3C288 has been omitted from our calculations of the %average jet power.  As a consequence, we may be underestimating the %mean jet power due to these rare but powerful sources.  \\subsection{AGN heating}\\label{sec:agnheating}   \\begin{deluxetable}{cc|cc|cc} \\tabletypesize{\\scriptsize} \\tablewidth{0pc} \\tablecolumns{6} \\tablecaption{Energy per particle\\label{table:hpb}} \\tablehead{ \\multicolumn{2}{c}{}                                                                       & \\multicolumn{2}{c}{$<$250\\,kpc}                                 & \\multicolumn{2}{c}{$<$500\\,kpc} \\\\ \\colhead{N$_{\\rm cl}$} & \\colhead{${\\bar{L}_{\\rm X,bol}}$}  & \\colhead{N$_{\\rm det}$} & \\colhead{heating}     & \\colhead{N$_{\\rm det}$} & \\colhead{heating}     \\\\ \\colhead{}                        & \\colhead{$10^{44}$\\ergs}              &  \\colhead{}                         & \\colhead{keV\\,par$^{-1}$}     & \\colhead{}                          & \\colhead{keV\\,par$^{-1}$}  } \\startdata 65 & 0.47        &   19 &   $0.23\\pm0.12$     & 56 &        $0.39\\pm0.17$ \\\\ 57 & $1.73$   &   19 &  $0.12\\pm0.06$      & 42 &        $0.27\\pm0.13$ \\\\ 47 & $6.97$   &   16 &  $0.17\\pm0.10$      & 36 &        $0.27\\pm0.14$ \\\\ \\cline{1-6} \\vspace{-2.5mm} \\\\ 169 & $2.71$ &  54  &  $0.17\\pm0.08$     & 134&       $0.30\\pm0.14$ \\enddata \\end{deluxetable}   From the jet power, we can estimate the mean total energy per particle injected by the radio sources (Table~\\ref{table:hpb}) \\begin{equation} {\\rm E}_{\\rm jet}=\\meanpcav\\,t_{\\rm z}{\\mu}m_{\\rm p}/\\bar{M}_{\\rm gas}, \\end{equation} where ${\\mu}=0.59$ is the mean molecular weight and $m_{\\rm p}$ is the proton mass. $\\meanpcav$ and $\\bar{M}_{\\rm gas}$ are the mean jet power and gas mass, respectively, for the clusters\\footnote{The mean jet power we estimated is defined within an aperture projected on the sky, but the gas mass here is defined in three dimension. From Fig.~\\ref{fig:radialdistr} and discussions in \\S\\ref{sec:distribution_radio}, the radio sources in clusters are dominated by central sources. Therefore, the difference between projected and three dimensional mean jet power should not be significant, especially for small aperture like 250\\,kpc.}.  The gas mass is calculated from the X-ray luminosity assuming the $L_{\\rm X}$-$M_{500}$ relation in \\citet{vikhlinin09} and a gas mass fraction of $0.12$.  We excluded clusters beyond $z = 0.6$ because their numbers are small.  The duration of energy injection, $t_z$, is estimated as the time interval from $z=0.6$ to $0.1$ (4.4 billion years).  As might be expected from the discussion of Fig.~\\ref{fig:cavitypower_lum_lum}, ${\\rm E_{\\rm jet}}$ decreases with X-ray luminosity, although the trend is weaker due to the luminosity dependence of the $L_{\\rm X}$-$M_{500}$ relation.  AGN heating from within the core alone can reach about $0.2$\\,keV/particle for poor clusters with typical X-ray luminosities of $5\\times10^{43}$\\ergs.  The number increases to $\\sim 0.4$\\,keV/particle if the aperture increases to 500\\,kpc. We have also combined the data for the three X-ray luminosity bins, giving the average AGN energy input per particle for the entire sample in the last column of the table.   We expect the average energy injected by radio sources to be underestimated somewhat due to the radio luminosity cut, particularly for the high redshift clusters.  We examine this effect by calculating the contribution from the weaker radio sources (${\\rm P}_{1.4}<10^{24}$\\,W/Hz) in clusters lying between redshift 0.1 and 0.3. We further assume that the radio luminosity function does not evolve with redshift.  Including the contribution of the weaker sources, the average energy injected by radio sources into clusters at $0.3<{\\rm z}<0.6$ is boosted by $7\\%$ within 250\\,kpc radius and by $11\\%$ within 500\\,kpc.  This correction ignores radio sources lying below our 3\\,mJy radio flux limit.  Lowering the limit to $2$\\,mJy and weighting to offset to the flux-dependent completeness of the NVSS survey \\citep{nvss}, the increase in the average injected energy is negligible ($\\le 2\\%$). Since the dependence of jet power on radio power according to Cavagnolo's calibration is steeper than found previously by \\citet{birzan08}, the contribution of weak source is expected to be less important than was found, for example, by \\citet{hart09}.   We note that the numbers in the table are calculated using the mass of gas within R$_{500}$. The radio sources considered here deposit most of their energy directly into the cluster cores, but current ignorance of the significant heating processes means that we do not know where the bulk of the energy will ultimately reside.  If it did remain within the cluster cores, the energy deposited per unit mass would be approximately $\\sim5$ times\\footnote{The factor is estimated from the ratio of mass within  R$_{500}$ and 250\\,kpc radius assuming a simple density profile $\\rho \\propto r^{-2}$.}) larger than the values given in the table, with a commensurate increase in its local impact. However, our primary objective is to demonstrate that the energy input is significant, even if the energy ends up being spread uniformly throughout much of the cluster atmospheres.   We have limited our jet power calculation in the redshift interval of the sample  ($0.1<\\rm z<0.6$). In principle, the energy injected by radio sources should be traced back to the time when BCGs formed \\citep[$z\\sim2.0$, e.g.][]{vandokkum01}.  Ignoring AGN evolution, the energy injected per particle since $z=2$ would be at least twice the values listed in Table~\\ref{table:hpb}.  This is a conservative estimate, since the fraction of powerful AGN increases with redshift \\citep[e.g.][]{galametz09,martini09} and the contribution of AGN heating is expected to be significantly greater in the past.  Furthermore, we should take into account the evolution of clusters over this time scale, so that we cannot simply scale the values in the first three rows of Table~\\ref{table:hpb}, which are derived for fixed X-ray luminosity ranges.  Na\\\"ively, if we assume that the mass function of 400SD clusters is representative of the ``integrated\" mass function of clusters since $z=2$, we can extrapolate the numbers in the last row of Table~\\ref{table:hpb} to provide a rough estimate of the minimum AGN energy injected per particle since $z=2$. The values we find are $0.4$\\,keV/particle from AGN within 250\\,kpc.  \\citet{wu00} found that the minimum excess energy required to break self-similarity is $\\simeq 1$\\,keV/particle.  Therefore, our results suggest that the continuous input of energy from AGN activity at the current rate over the ages of clusters would provide a significant fraction (40\\%) of the excess entropy (heat) found in the hot atmospheres of clusters.  It also implies that AGN heating may be a factor leading to the dearth of strong cooling flows in distant clusters \\citep[e.g.][]{vikhlinin07,samuele11}.  It is unclear whether this heating is being supplied by a self-regulated feedback loop.      We have correlated the positions  of clusters from the 400SD cluster survey lying between redshifts 0.1 and 0.6  against radio sources  with  fluxes above 3\\,mJy from the NVSS. Radio sources are detected within 250\\,kpc of the center for 30\\% of the clusters and within 1Mpc for 50 to 80\\%, depending on redshift. The first value agrees with the fraction of radio sources in BCGs found by \\citet{best07a} for SDSS clusters, but  lies well below the nearly 100\\% detection rate in strong cooling flows \\citep[e.g.][]{burns90,edwards07, dunn06b, mittal09, cavagnolo09}. We found that radio power of radio sources in the 400SD clusters does not correlate with the cluster X-ray luminosity. For comparison, \\citet{mittal09} and \\citet{sun09} found strong correlations between AGN radio power of BCG host galaxies and the X-ray luminosities of their bright cooling cores, and the correlation disappeared in systems with weak or absent cooling cores.    Both of these facts support our conclusion that the 400SD sample is composed primarily of weak or non-cooling core clusters.  We have found no significant correlation between the incidence of radio sources in clusters and the X-ray luminosity. Since the X-ray luminosity of a cluster correlates with its richness, this implies that richer clusters have fewer radio AGN per galaxy. In addition, the density profile of radio sources in clusters peak at the cluster center, implying that BCGs are the dominant radio galaxies in clusters. Apart from the BCG, the rich cluster environment is hostile to the formation of radio AGN, so that the number of radio AGNs does not increase with the number of galaxies.  We have found only modest correlation between radio power and incidence of radio detection in 400SD clusters with redshift. Although it is difficult to completely rule out bias in a flux-limited sample, our results are qualitatively consistent with redshift evolution of the radio AGN detection rate in clusters by e.g. \\citet{galametz09}.  We have estimated the average jet power and the rate of AGN heating in clusters using the scaling relation of Eqn.~\\ref{eqn:high}. The average jet power is about $2\\times 10^{44}$\\ergs for radio AGNs within 250\\,kpc radius of the 400SD clusters and about $4\\times 10^{44}$\\ergs for sources within 500\\,kpc radius. We found that the average jet power and AGN heating rate do not correlate with total X-ray luminosity of the 400SD clusters. Therefore, the heating rate per particle will be larger in less massive systems. The AGN heating within the core of clusters can reach 0.2\\,keV/particle for poor cluster with typical X-ray luminosities of $5\\times 10^{43}$\\ergs. These numbers are calculated within the redshift range of our sample, i.e. $0.1<\\rm z<0.6$. If we extrapolate the result to redshift $\\rm z=2$ ignoring AGN evolution, the integrated AGN heating per particle will increase by a factor of 2.  If the heated gas is unable to cool quickly enough, the entropy of their hot atmospheres  will rise above the values expected from gravitational heating alone.  In fact, we found in \\S\\ref{sec:agnheating} that the amount of AGN heating of the hot atmospheres is a significant fraction ($40\\%$) of the heating required to ``preheat'' clusters \\citep[][]{kaiser91,wu00,mccarthy02}.  This so-called preheating phase is thought to occur during the epoch of galaxy formation at redshifts of 3 and beyond \\citep[][]{kaiser91}.  Our results suggest that the  AGN heating of cluster atmospheres long after the epochs of galaxy and cluster formation, and throughout the formation history of clusters are significant.  Thus, the heating that apparently broke the self-similarity of cluster scaling relations appears to have occurred continuously, and not necessarily at a single epoch."
    },
    "1107/1107.1004_arXiv.txt": {
        "abstract": "The tachocline is believed to be the region where the solar dynamo operates. With over a solar cycle's worth of data available from the MDI and GONG instruments, we are in a position to investigate not merely the average structure of the solar tachocline, but also its time variations. We determine the properties of the tachocline as a function of time by fitting a two-dimensional model that takes latitudinal variations of the tachocline properties into account. We confirm that if we consider central position of the tachocline, it is prolate. Our results show that the tachocline is thicker at higher latitudes than the equator, making the overall shape of the tachocline more complex.  Of the tachocline properties examined, the transition of the rotation rate across the tachocline, and to some extent the position of the tachocline, show some temporal variations. ",
        "introduction": "\\label{sec:intro}  Inversions of helioseismic data have shown that the solar convection zone rotates differentially, but that the radiative interior has an almost solid-body like rotation (see Schou et al.~1998 and references therein). The transition from differential rotation to solid body rotation occurs over a very narrow region, that is called the tachocline (Spiegel \\& Zahn 1992). The tachocline is believed to be the seat of the solar dynamo, and hence is expected to affect, and be affected by, the solar dynamo.  The poor resolution of the inversion techniques, particularly at high latitudes, makes it difficult to infer the properties of the tachocline from normal inversions for the solar rotation rate. One consequence of the poor resolution is that the transition appears to occur over a large radial distance. As a result, tachocline studies usually involve forward modelling techniques, with the parameters of the models determined by fitting helioseismic data (e.g., Kosovichev 1996; Basu 1997; Antia et al.~1998; Charbonneau et al.~1999; Basu \\& Antia 2001, 2003). Early investigations have shown that the tachocline is prolate. Also clear (even from inversions) is that the `jump' in the rotation rate across the tachocline is latitude dependent. At the equator, the rotation rate changes from a higher value to a lower value when moving from the convection zone to the radiative interior through the tachocline, there is almost no change around a latitude of about $30^\\circ$, while at higher latitudes the rotation rates shifts from a lower value to a higher one.  What is not completely clear from  earlier work however, is whether the tachocline thickness, i.e., the radial distance over which the rotation rate changes, is larger at higher latitudes than the equator. Although, the early results found that the thickness does increase with latitude, the statistical significance of this increase was not clear.  The bulk of the solar convection zone shows clear, periodic, changes in the rotation rate in the form of zonal flows (Vorontsov et al.~2002; Basu \\& Antia 2003; Howe et al.~2005). It is, however, not clear if the tachocline too changes with time. Early studies by Basu \\& Antia~(2001; 2003) using data available for a limited part of the solar cycle 23 during its ascending phase did not find any significant temporal variation. However, with more than a solar cycle worth of data available now, we revisit the question of tachocline variations. We also examine whether the tachocline structure was different between the solar minimum preceding cycle 23 and the exceptionally long and deep minimum that preceded cycle 24. We also use the increased amount of data  to determine the average properties of the tachocline more accurately. ",
        "conclusions": "\\label{sec:disc}  We have used helioseismic data spanning cycle 23 and beyond to study the properties of the solar tachocline. We confirm that the center of the tachocline is prolate in shape. The results show unequivocally that the thickness of the tachocline increases with increasing latitude, making the overall shape of the latitude more complex --- while the outer boundary is prolate, the inner boundary is close to spherical or perhaps even a little oblate. This appears to be consistent with the fact that the base of the convection zone is almost spherical.  The jump across the tachocline is the only parameter that shows a significant change with solar activity. Other parameters also appear to show some change with time, however the tachocline properties obtained with GONG data and those with MDI do not always show consistent behavior. If the oscillatory temporal variations of tachocline properties are confirmed, it would imply that the solar zonal flow (which is the temporally varying component of solar rotation) penetrates to the base of the convection zone."
    },
    "1107/1107.4767_arXiv.txt": {
        "abstract": "The Standard Model Higgs boson with large nonminimal coupling to the gravitational curvature can drive cosmological inflation. We study this type of inflationary scenario % in the context of supersymmetric grand unification and point out that it is naturally implemented in the {\\em minimal} supersymmetric $SU(5)$ model, and hence virtually in any GUT models. It is shown that with an appropriate K\\\"{a}hler potential the inflaton trajectory settles down to the Standard Model vacuum at the end of the slow roll. The predicted cosmological parameters are also consistent with the 7-year WMAP data. ",
        "introduction": "Recently the idea that the Standard Model (SM) Higgs field may be identified with an inflaton field, has attracted much attention \\cite{ BS,Barvinsky:2008ia,DeSimone:2008ei,BKKSS:2009,Barbon:2009ya,Bezrukov:2010jz,Burgess,Hertzberg:2010dc,LernerMcDonald}. The major r\\^{o}le is played by the nonminimal coupling to gravity, which renders the Higgs mass to be within the range of $126 - 194$ GeV \\cite{BS,Barvinsky:2008ia,DeSimone:2008ei,BKKSS:2009}, while keeping the amplitude of the primordial curvature perturbation at the scale of $\\sim 10^{-5}$. The idea of inflation by nonminimally coupled inflaton field itself is certainly not new \\cite{nonminimal}. Nevertheless, the striking agreement with the present-day cosmological data, combined with the minimalistic nature of the model, makes this type of scenario very attractive. The predicted mass range of the Higgs particle is also interesting for the physics of the Large Hadron Collider.   The Higgs potential in the SM is unstable against quantum corrections (the hierarchy problem) and it therefore is reasonable to reconsider Higgs inflation in supersymmetric theory % \\cite{Einhorn:2009bh,FKLMvP}. It is shown in \\cite{Einhorn:2009bh} that Higgs inflation cannot be implemented within the minimal supersymmetric Standard Model (MSSM), as the field content of the latter is too restrictive. Instead, with an extra field (i.e. in the next-to-minimal supersymmetric Standard Model, NMSSM) a sensible scenario of Higgs inflation is found to be possible. The NMSSM model has tachyonic instability in the direction of the extra field, but this can be cured by considering a noncanonical K\\\"{a}hler potential \\cite{FKLMvP}.  In this paper we discuss the possibility of Higgs inflation in supersymmetric grand unified theory (GUT). There are several reasons to motivate this study. One obvious reason is that the energy scale of inflation is typically above the grand unification scale, and it is unnatural to suppose that the SM Lagrangian is valid all the way up to the scale of inflation; as the GUT scale destabilises the electroweak scale without supersymmetry, it seems that supersymmetric GUT is an appropriate theory to start with. Another reason is the puzzling necessity of the extra field besides the MSSM fields for successful Higgs inflation, as alluded to above; going beyond the MSSM is somewhat against the minimalistic guiding principle of the original Higgs inflation, and as the NMSSM is structurally similar to the $SU(5)$ GUT model, it seems natural to conjecture that the $SU(5)$ GUT, rather than the NMSSM, may be a more appropriate minimal supersymmetric theory that accommodates Higgs inflation. Obvious questions are then whether it is possible to obtain enough inflation (e-folding) somewhere between the Planck scale and the GUT scale, and if so whether the prediction of the cosmological parameters is consistent with the present observation. We shall address these issues below, and find that a viable Higgs inflationary scenario nicely fits into the minimal $SU(5)$ model. We shall employ supergravity embedding of GUT \\cite{sGUT}, since the nonminimal coupling of the Higgs field to gravity naturally arises in that framework.  \\begin{figure*} \\begin{eqnarray*} \\begin{array}{ccccc} \\includegraphics[width=60mm]{3d.eps}% &~~~& \\includegraphics[width=40mm]{contour.eps}% &~~~& \\includegraphics[width=60mm]{pot.eps}% \\end{array} \\end{eqnarray*} \\caption{\\label{fig:VE} The scalar potential $V_{\\rm E}$ in the Einstein frame (left), the inflaton trajectory in the contour plot of the same potential (middle), and the minima of the scalar potential $V(s(h),h)$ plotted against $h$ (right). In the middle panel the thick red curve is the inflaton trajectory. We have chosen $\\rho=0.5$, $\\lambda=0.5$, $\\omega=-100$, $\\zeta=10000$. The nonminimal coupling $\\gamma=1.86\\times 10^4$ is fixed by the amplitude of the curvature perturbation, evaluated for e-folding $N_{\\rm e}=60$. } \\end{figure*} ",
        "conclusions": "\\label{sec:Discussion} In this paper we have discussed Higgs inflation in supersymmetric GUT, taking the minimal $SU(5)$ model as a concrete example. In the early days the proposals of cosmological inflation were made for the Higgs field in the GUT models \\cite{inflation}. % It is intriguing to see that the prediction based on the simplest GUT, with the help of nonminimal coupling to gravity, is in perfect fit with today's observational constraints. %   The nonminimal coupling is consistent with the symmetries of general relativity and the SM, and it naturally arises in quantum field theory in curved spacetime \\cite{Birrell:1982ix}. % The value of the coupling $\\xi\\sim 10^4$, however, is rather large. This is a generic feature of Higgs inflation, since successful slow-roll requires $h^2\\lesssim M_{\\rm P}^2\\lesssim\\xi h^2$ \\cite{BS}. It has been argued that such large nonminimal coupling could violate the unitarity bound, since the cut-off scale evaluated as $M_{\\rm P}/\\xi$ is considerably lower than the Planck scale \\cite{Barbon:2009ya,Burgess,LernerMcDonald,Hertzberg:2010dc}. Others contend that such a criticism is not valid, arguing that at large field values $\\gtrsim M_{\\rm P}/\\xi$ the cut-off scale is actually field-dependent \\cite{BKKSS:2009,Bezrukov:2010jz,FKLMvP}. The large nonminimal coupling is, at any rate, a key feature of the Higgs inflation and it is certainly worthwhile understanding possible dangers arising from this. Another type of criticism concerns the quantum stability of the classical potential. This problem was studied using renormalisation group (RG) analysis \\cite{BKKSS:2009,Barvinsky:2008ia,DeSimone:2008ei}, and the effects of renormalisation are found to be small except for some extreme values of parameters. We have also performed RG analysis in our GUT model and verified that the effects are small (less than $3\\%$ for $r$, less than $2\\%$ for $\\xi$, and less than $0.1\\%$ for $n_s$). This is expected, since inflation takes place in a narrower energy range of $10^{16} - 10^{18}$ GeV and the RG effects should be smaller than the SM case.   A closer look at the potential $V_{\\rm E}$ shows that its minimum is at a small negative value, $\\sim -2 \\times 10^{-16} M_{\\rm P}^4$, for our parameter choices. This is offset by a contribution from the supersymmetry breaking sector and the scenario does not suffer from the cosmological constant problem. In our scenario the energy scale of inflation is in the GUT scale and the Higgs fields are directly coupled to the SM particles. This indicates that the reheating temperature is high, typically from the intermediate to the GUT scale. It would be interesting to discuss further phenomenological implications, such as the gravitino problem and baryogenesis.   In this paper we considered a single-field Higgs inflation model appropriate for our parameter choice $\\zeta=10000$, $\\omega=-100$ of the K\\\"{a}hler potential. These values are not too exotic, as $\\langle\\Phi\\rangle$ is still very close to $1$ and the Planck scale after inflation is nearly $M_{\\rm P}$. For smaller values of $\\zeta$ and $|\\omega|$, the displacement of $s$ during inflation becomes large. This leads to two-field inflation, which is also of interest, in particular, due to possible generation of detectable large non-Gaussianity. Supersymmetric models of Higgs inflation necessarily involve multiple fields \\cite{Einhorn:2009bh}. The engendered isocurvature mode can, in principle, distinguish various models of Higgs inflation.   Finally, the scenario can also be extended to other GUT models whose gauge group contains $SU(5)$ as a subgroup. When the Higgs multiplets of the GUT model contain ${\\bf 5}$, ${\\bf \\bar{5}}$ and ${\\bf 24}$ of the minimal $SU(5)$ GUT, a superpotential like (\\ref{eqn:Super}) can be introduced. Then a viable model of Higgs inflation is implemented, as described in this Letter. One such simple example is the $SO(10)$ GUT with Higgs multiplets in ${\\bf 10}$ and ${\\bf 54}$ representations.   \\subsection*"
    },
    "1107/1107.4551_arXiv.txt": {
        "abstract": "It has recently been argued by Copeland {\\it et al.} that in eleven dimensions two orbifold planes can collide and bounce in a regular way, even when the bulk metric is perturbed away from Milne spacetime to a Kasner solution. In this paper, we point out that as a consequence the global ``phoenix'' structure of the cyclic universe is significantly enriched. Spatially separated regions, with different density fluctuation amplitudes as well as different non-gaussian characteristics, are all physically realized. Those regions containing by far the most structure are specified by a fluctuation amplitude of $Q \\sim 10^{-4.5}$ and local non-gaussianity parameters $f_{NL} \\sim {\\cal O}(\\pm 10)$ and $g_{NL} \\sim {\\cal O}(-10^3),$ in agreement with current observations. ",
        "introduction": "In cosmology, it is tempting to declare that we should only ever think about what happened in our past lightcone, as everything outside of it is inaccessible to observations. However, several well-motivated models of the early universe predict the existence of regions of the universe, with a whole range of physical properties, that are far separated in space and/or in time from our observable region of the universe. Hence it may be important to understand the characteristics of these far removed regions, at least on a theoretical level. It certainly seems of interest to explore all the consequences of a given cosmological model, whether these appear to be immediately observable or not; this is crucial in assessing the logic and internal consistency of the model in question, and hence crucial in determining whether a given cosmological model satisfies us intellectually. Such an understanding has concrete implications, as the perception of which properties of our region of the universe we consider to be determined by mathematical necessity rather than historical accident guides us in devising future research avenues.  The question of additional universes is particularly acute for eternal inflation \\cite{Guth:2007ng}, since the number and variety of universes are both infinite\\footnote{We will occasionally use the expressions ``universe'' and ``region of the universe'' interchangeably, when this does not seem confusing.}. Hence all possible universes are produced an infinite number of times, so that, taken by itself, the model of eternal inflation is not predictive. In order to make it predictive one needs to add a measure, which regulates the infinities in question and leads to definite predictions. Unfortunately it turns out that all predictions depend sensitively on the measure/regulator used \\cite{Freivogel:2011eg} -- in other words, unlike the situation in QED for example, all predictions are regulator-dependent. Moreover, some of the simplest measures turn out to be in blatant conflict with observations. Hence we are certainly missing something important, and one may even conjecture that eternal inflation, as currently envisaged, is unphysical and does not occur.  In this paper, we will examine the question of predictivity in an alternative cosmological model, namely the cyclic theory of the universe proposed by Steinhardt and Turok \\cite{Steinhardt:2001st}. In doing so, we will extend existing studies of the global structure of the cyclic universe, previously described as the {\\it phoenix universe} \\cite{Lehners:2008qe,Lehners:2009eg}. Building on recent results by Copeland, Niz and Turok regarding brane collisions \\cite{Copeland:2010yr}, we will show that the cyclic universe also leads to vastly separated regions with distinct physical properties. Hence one may be worried that a similar ambiguity as the one inherent to eternal inflation will come to plague the model. However, as we will describe, the cyclic model predicts strong correlations between a large number of physical observables in a given region, such as its spatial flatness, the amount of structure it contains and the non-gaussian statistics of its density perturbations. These correlations are entirely determined by the dynamics of the model, and thus they lead to clear, testable, predictions. For the specific model described below, it turns out that those regions that are flat and contain the most structure are characterized by a fluctuation amplitude of $Q \\sim 10^{-4.5}$ and local non-gaussianity parameters $f_{NL} \\sim {\\cal O}(\\pm 10)$ and $g_{NL} \\sim {\\cal O}(-10^3).$ Thus the model is both in agreement with current observations, and falsifiable by upcoming CMB experiments.  Before proceeding, we should clarify that the cyclic model contains a number of open questions. The most important such open question has to do with the bounce that connects the ekpyrotic contracting phase with the phase of expansion. This big crunch/big bang transition has been modeled by a collision of orbifold branes in M-theory \\cite{Khoury:2001wf}. To date, it has only been possible to analyze such a collision at the semi-classical level \\cite{Turok:2004gb,Copeland:2010yr}, and not at the full quantum level. In the present work, we take the semi-classical results seriously, and explore their consequences. Thus, the present work is only applicable insofar as the semi-classical calculations accurately reflect the physics of the bounce. But, to the extent that our analysis yields promising results, it also motivates further work on the quantum treatment of the crunch/bang transition.  The plan of the paper is as follows: we first provide a brief review of the cyclic universe, focussing on those aspects that are most relevant to the present paper. Section \\ref{SectionPhoenix} then discusses the global structure of the cyclic universe, reviewing the picture as it was understood up to now, and extending it by using the results of recent semi-classical calculations of colliding orbifolds. The new global phoenix structure of the cyclic universe, together with its implications, is discussed. Conclusions and further directions are to be found in the discussions section. ",
        "conclusions": "The best understood version of the cyclic universe involves an ekpyrotic phase during which the potential is unstable. Because of this instability, nearly scale-invariant isocurvature perturbations are generated, which get converted into curvature perturbations as the trajectory in field space turns. This turn can happen in essentially two ways, either during the ekpyrotic phase or during the subsequent kinetic phase. It was thought up to now that these possibilities correspond to different cosmological models, depending on how the initial conditions are arranged. Here, we have shown that in the cyclic universe, where each cycle generates the ``initial conditions'' for the next one, all of these possibilities are physically realized in parallel.  This leads to a global structure of the phoenix type, where large, flat regions now come in a variety of physical characteristics, but, interestingly, with strong correlations between a number of cosmological observables. The regions containing by far the most structure have a primordial density fluctuation amplitude of $Q_{\\mathrm{max}} \\approx 10^{-4.5},$ with local non-gaussian corrections specified by $f_{NL} \\sim {\\cal O}(\\pm 10)$ and $g_{NL} \\sim {\\cal O}(-1000).$ There are far more numerous regions with a continuous range of ever smaller density perturbation amplitudes (starting at about $Q_{\\mathrm{max}}/2$) and local non-gaussianity parameters $f_{NL} \\sim {\\cal O}(- 10)$ and $g_{NL} \\sim {\\cal O}(+1000).$ Even when counted together, these regions contain far fewer large amplitude fluctuations though than the regions having the maximal possible density fluctuation amplitude, and thus most galaxies are to be found in the $Q_{\\mathrm{max}}$ regions. Also, as is typical in cyclic models, none of these regions are expected to contain large amplitude primordial gravitational waves. Finally, there are regions that are in a gravitationally collapsed state. Those are regions where the ekpyrotic phase ended prematurely, and which (we conjecture) underwent a high-curvature brane collision from which they emerged with an overproduction of matter, causing them to re-collapse rapidly.  One can summarize the essential steps that lead to the global phoenix structure as follows: diversify, filter, amplify. The unstable potential diversifies the possible physical properties that the model produces, the brane collision filters out those regions that are sufficiently flat, and the radiation, matter and dark energy phases amplify those regions that made it through the bounce in good shape. This sequence of steps contributes to making the model highly predictive. Another aspect enhancing its predictivity is the fact that during the longest smoothing phase, \\ie during the dark energy phase, the Hubble rate is lower than during the principal phases of structure formation (\\ie the ekpyrotic, radiation and matter phases). This implies that even large quantum fluctuations during the dark energy phase do not change the structure of the model in an essential way (in contrast to eternal inflation). The only consequence of such fluctuations are that the corresponding regions expand a little more, and enter their next cycle a little later than the surrounding regions. A more quantitative analysis of this particular point will be left for future work.  From a model-building perspective, the new phoenix universe presented here has two advantages over the ``old'' version. First, it shows that no fine-tuning is needed regarding the orientation of the unstable ridge of the ekpyrotic potential. And secondly, the constraint on the number of e-folds of dark energy needed to sustain the cyclic universe is weakened. In fact, in the new picture, it is likely that our predecessor-universe was a region with low $Q$ and very little structure.  In this framework, there are many opportunities for further study: one may wonder for example how sharp the boundaries between the different regions in the phoenix universe are, and whether these boundaries lead to any interesting effects. Evidently, it would be desirable to have a fully quantum calculation of the bounce, but such a treatment may still be quite far into the future. In the meantime, it would certainly be worthwhile to study more detailed models of the bounce, and perhaps investigate under what circumstances physical quantities can change. An interesting possibility is that the bounce may be a strong filter for many other physical quantities too, and thus may select a whole range of features of our universe that otherwise might appear to be merely historical accidents."
    },
    "1107/1107.3147_arXiv.txt": {
        "abstract": "We utilize for the first time HST ACS imaging to examine the structural properties of galaxies in the rest-frame $U-V$ versus $V-J$ diagram (i.e., the $UVJ$ diagram) using a sample at \\zwindowwide\\ that reaches a low stellar mass limit (\\mcutrangelog).  The use of the $UVJ$ diagram as a tool to distinguish quiescent galaxies from star forming galaxies (SFGs) is becoming more common due to its ability to separate red quiescent galaxies from reddened SFGs.  Quiescent galaxies occupy a small and distinct region of $UVJ$ color space and we find most of them to have concentrated profiles with high S\\'{e}rsic indices ($n>2.5$) and smooth structure characteristic of early-type systems.  SFGs populate a broad, but well-defined sequence of $UVJ$ colors and are comprised of objects with a mix of S\\'{e}rsic indices.  Interestingly, most $UVJ$-selected SFGs with high S\\'{e}rsic indices also display structure due to dust and star formation typical of the $n<2.5$ SFGs and late-type systems.  Finally, we find that the position of a SFG on the sequence of $UVJ$ colors is determined to a large degree by the mass of the galaxy and its inclination.  Systems that are closer to edge-on generally display redder colors and lower \\oii\\ luminosity per unit mass as a consequence of the reddening due to dust within the disks.  We conclude that the two main features seen in $UVJ$ color space correspond closely to the traditional morphological classes of early and late-type galaxies. ",
        "introduction": "\\begin{figure*} \\epsscale{1} \\plotone{f1.eps} \\caption{($a$) Rest-frame $U-V$ vs. $V-J$ colors for \\nmcut\\ galaxies with mass \\mcutrangelog\\ at \\zwindowwide.  The typical $\\pm 1$-$\\sigma$ statistical uncertainty in the rest-frame colors is indicated in the bottom right.  The \\citet{williams2009} boundary (black wedge) separates quiescent galaxies (top left) from SFGs (bottom right).  The red line indicates the evolution of a BC03 solar metallicity single burst, while the blue line represents a constant SFR model (CSF).  The diamonds on each model track represent time steps of 0.1, 0.5, 1, 2, 3, and 5 Gyr (from bottom to top).  The arrow indicates the reddening vector.  ($b$) Gray scale representation of the density of points in panel ($a$).  The concentration of quiescent galaxies in the $UVJ$ plane is distinct from the broader distribution of SFGs.  The dashed line runs through the $t=5$~Gyr CSF time step and is parallel to the reddening vector.  We examine axis ratios and \\oii\\ luminosities along this line, \\vjproj, in Figure~\\ref{fig_boa_VJproj} as we hypothesize that reddening is a major factor in explaining the sequence of $UVJ$ colors for SFGs.} \\label{fig_UVJ_letter} \\end{figure*}   In the absence of morphological information to distinguish early and late-type galaxies, the use of UV-optical colors has emerged as a commonly employed proxy in studies of galaxy evolution.  This practice is bolstered by the color bimodality \\citep[e.g.,][]{strateva2001,baldry2004}.  For example, galaxies are often divided in a color-magnitude or color-mass diagram into red and blue, and the evolution in the luminosity or mass function of each color type tracked over cosmic time \\citep[e.g.,][and several others]{bell2004b,faber2007,brown2007}.  The underlying assumption in the method above is that the red galaxies represent quiescent systems with little or no ongoing star-formation (SF).  However, it is well known that a reddened star forming galaxy (SFG) can display colors coincident with those on the red sequence \\citep[e.g.,][]{maller2009}.  Selecting red galaxies based on a single rest-frame color, in combination with a magnitude or stellar mass, therefore results in a selection of both quiescent galaxies {\\em and} reddened SFGs.  This presents a predicament for studies that use such a selection in order to quantify the evolution in the number of passive systems.  In this Letter, we highlight the utility of a color-color diagram, specifically rest-frame $U-V$ vs. $V-J$ (hereafter referred to as the $UVJ$ diagram), in distinguishing quiescent galaxies from SFGs, including those SFGs that are heavily reddened.  The $UVJ$ diagram is gaining increasing visibility in studies of galaxy evolution \\citep[e.g.,][]{labbe2005,wuyts2007,williams2009,brammer2009,williams2010,whitaker2010,patel2011,quadri2011,brammer2011}.  In \\citet{patel2011} we hinted at the possibility of the $UVJ$ selection of quiescent and SFGs being used to also distinguish morphologically early and late-type systems.  Here we introduce for the first time an analysis based on {\\em HST} ACS imaging and examine the structural properties of galaxies in $UVJ$ color space in order to address the prospects of using two rest-frame colors in classifying both the recent SFHs and morphologies of galaxies.  We use a stellar mass-limited sample of galaxies at \\zwindowwide\\ to carry out our analysis.  As is well known, many galaxy properties correlate with stellar mass and it is therefore crucial that one control for stellar mass in interpreting results.  The redshift range examined here provides a suitable case study for a sample of galaxies in the distant universe.  The rest-frame $J$-band magnitudes at these redshifts represent observed $K_s$, which is typically the reddest bandpass accessible from the ground for deep wide-field galaxy surveys.  We assume a cosmology with $H_0=70$~\\kmps~\\pmpc, $\\Omega_M = 0.30$, and $\\Omega_{\\Lambda} = 0.70$.  Stellar masses are based on a Chabrier IMF \\citep{chabrier2003}.  All magnitudes are reported in the AB system. ",
        "conclusions": "In the rest-frame $UVJ$ diagram, quiescent galaxies occupy a concentrated region of color space that is distinct from the sequence of $UVJ$ colors of SFGs, which extends onto the red sequence.  Utilizing {\\em HST} ACS imaging for the first time to study the structural parameters of galaxies in the $UVJ$ diagram at \\zwindowwide, we find that the commonly employed $UVJ$ selection of quiescent and SFGs corresponds closely to the distinction between traditional morphological classes of early and late-type galaxies.  Meanwhile, the position of a SFG along the extended sequence of $UVJ$ colors is determined to a large degree by its stellar mass and viewing inclination.  Systems that are closer to edge-on generally display redder $UVJ$ colors and have lower \\oii\\ luminosity per unit mass as a consequence of the reddening from dust within the inclined disks.  Thus, most of the $UVJ$-selected SFGs on the optical red sequence are highly inclined spirals.  The use of only two rest-frame colors in classifying both the recent SFHs and morphologies of galaxies will be invaluable for future studies of galaxy evolution."
    },
    "1107/1107.3153_arXiv.txt": {
        "abstract": "We present 880\\,$\\mu$m Submillimeter Array observations of the debris disks around the young solar analogue HD~107146 and the multiple-planet host star HR~8799, at an angular resolution of 3'' and 6'', respectively.  We spatially resolve the inner edge of the disk around HR~8799 for the first time.  While the data are not sensitive enough (with rms noise of 1\\,mJy) to constrain the system geometry, we demonstrate that a model by \\citet{su09} based on the spectral energy distribution (SED) with an inner radius of 150\\,AU predicts well the spatially resolved data.  Furthermore, by modeling simultaneously the SED and visibilities, we demonstrate that the dust is distributed in a broad (of order 100\\,AU) annulus rather than a narrow ring. We also model the observed SED and visibilities for the HD~107146 debris disk and generate a model of the dust emission that extends in a broad band between 50 and 170\\,AU from the star.  We perform an {\\it a posteriori} comparison with existing 1.3\\,mm CARMA observations and demonstrate that a smooth, axisymmetric model reproduces well all of the available millimeter-wavelength data. ",
        "introduction": "The tenuous, dusty debris disks orbiting main sequence stars are an important tracer of planetary system architectures and their dynamical evolution.  With typical ages of 10-100\\,Myr, debris disks are analogous to our own solar system during the period when giant and terrestrial planets were undergoing heavy bombardment by small bodies.  By this age, nearly all the primordial molecular gas has been dispersed from the protoplanetary disk, leaving a population of dust grains that should be removed by stellar radiation pressure on timescales much shorter than the age of the star.  The dust is therefore thought to be generated by the grinding together of planetesimals that formed from the primordial disk material.  Approximately 20\\% of nearby main sequence stars exhibit detectable infrared excesses indicative of debris disks, implying that planet formation at least to the stage of planetesimal growth is a common process \\citep[e.g.,][]{hab01,car05,bry06,moo10,mat11}.  However, despite their prevalence in the solar neighborhood, the relatively faint emission from orbiting dust grains makes imaging challenging, and only a handful of systems have thus far been spatially resolved.   Observations at millimeter wavelengths have proven particularly valuable for studies of system dynamics, as these longer wavelengths are most sensitive to larger dust grains with long resonant lifetimes that trace best the structure resulting from orbital resonances \\citep{wya06}. To date, there have been only three debris disks spatially resolved with interferometers: $\\beta$ Pictoris, HD~107146, and HD~32297 \\citep{wil10,cor09,man08}.  In order to better understand the dust distribution in these systems and contribute to the sample of spatially resolved debris disk systems, we have observed two nearby, relatively bright debris disks with the Submillimeter Array (SMA), adding a new wavelength to the suite of observations of HD~107146 and spatially resolving the disk around HR~8799 for the first time.  HD~107146 is a rare example of a young, nearby, close solar analogue, with spectral type G2V \\citep{jas78} and a distance of only 28.5\\,pc from the Sun \\citep{per97}.  Its age is estimated to lie between 80 and 200\\,Myr, although its location on the H-R diagram allows for an age as young as 30\\,Myr \\citep{wil04}.  It was identified as an ``excess dwarf'' on the basis of its {\\it IRAS} colors by \\citet{sil00} and its debris disk was detected and marginally resolved for the first time at a wavelength of 450\\,$\\mu$m by \\citet{wil04}.  It has since been spatially resolved in near-infrared scattered light \\citep{ard04,ert11} as well as 350\\,$\\mu$m, 1.3\\,mm and 3\\,mm dust continuum emission \\citep{car05,cor09}.  Imaging reveals a broad debris belt with maximum brightness near a radius of $\\sim$100\\,AU from the star, although detailed modeling of the millimeter data has not yet been carried out to determine its extent.  \\citet{naj05} place an upper limit of 1 on the gas-to-dust mass ratio, indicating that the disk is largely devoid of primordial molecular gas.  \\citet{cor09} found that their 1.3\\,mm image exhibits possible azimuthal brightness asymmetries, and suggested the presence of a planet that is shepherding the material into resonances.  Further observations of the large dust grain population are necessary to investigate that assertion.  HR~8799 is so far the only star to host a directly-imaged multiple-planet system.  Four planets have been observed with projected orbital radii ranging from $\\sim$15-80\\,AU \\citep{mar08,mar10}.  HR~8799 also hosts a debris disk, and was one of the twelve originally identified ``Vega-like'' stars based on its 60\\,$\\mu$m {\\it IRAS} flux \\citep{sad86}.  SED modeling suggests that there are at least two dust belts in the system, likely bracketing the radii of the innermost and outermost planets known: the inner belt has a temperature of $\\sim$150\\,K and the outer belt has a temperature of $\\sim$45\\,K \\citep{wil06,che09,su09}. Like all isolated young A stars \\citep[spectral type AV5, see e.g.][]{gra99}, its age is extremely difficult to determine, which is unfortunate because the masses of its planetary companions are estimated from hot-start models that depend sensitively on the assumed time since formation.  \\citet{mar08} estimate an age between 30 and 160\\,Myr based on several different indicators, which implies planet-mass companions \\citep[see also discussion in][]{hin10}. \\citet{moy10} favor an age as old as 1\\,Gyr based on an asteroseismological analysis, which would imply that the companions are in fact brown dwarfs, alothough an infrared color analysis by \\citet{cur11} argues against such a high mass for the companions.  The presence of a molecular gas filament coincident in space and recession velocity with HR~8799 \\citep{wil06,su09} may also point to a younger age.  Spatially resolving the disk is therefore crucial, because it can provide new constraints on the companion masses.  Both the age determination and the stability of the planetary system depend sensitively on the unknown viewing geometry \\citep{fab10}, which can be determined by spatially resolved observations of the disk, assuming coplanarity.  Spatially resolving the dust distribution, particularly at long wavelengths, also places a more direct dynamical constraint on the mass of the outermost planet independent of the age of the system \\citep[see, e.g.,][]{chi09}.  To investigate the spatial distribution of large dust grains and their implications for inferring the presence and properties of planets in the systems, we observed HD~107146 and HR~8799 with the SMA at a frequency of 880\\,$\\mu$m.  The observations are described in Section~\\ref{sec:obs} and the results are presented in Section~\\ref{sec:results}.  We generate a model of the disk around HD~107146 that is capable of reproducing both the SED and SMA visibilities and compare it with the previous 1.3\\,mm CARMA observations in Section~\\ref{sec:hd_analysis}.  We perform a simplified analysis on the lower signal-to-noise HR~8799 data in Section~\\ref{sec:hr_analysis}, comparing the data with the SED-based model from \\citet{su09} and investigating whether the data can distinguish between a broad and narrow dust ring.  We discuss the implications of these results for the properties of the systems in Section~\\ref{sec:discussion}, and summarize our conclusions in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have obtained the first spatially resolved 880\\,$\\mu$m observations of the dusty debris disks around the young solar analogue HD~107146 and the multiple-planet host star HR~8799 at a resolution of 3\" and 6\", respectively. The data reveal broad bands of emission with inner radii at tens of AU from the central star.  Models of axisymmetric dust annuli are capable of reproducing both the SED and 880\\,$\\mu$m visibility data, with no significant residuals that would indicate clumpiness due to shepherding of dust grains into orbital resonances with planets.  The outer debris belt around HR~8799 is spatially resolved for the first time.  While the spatially resolved data alone are not of sufficiently sensitive to constrain the location of the inner edge of the HR~8799 disk, the SED-based model of \\citet{su09} with an inner edge at 150\\,AU from the star is shown to be consistent with the SMA data.  A 150\\,AU inner radius serves roughly as the largest radius at which the debris belt may be plausibly truncated by the chaotic zone of HR~8799b, and more sensitive imaging better able to pinpoint the location of the inner edge of the debris belt would be extremely valuable.  The 880\\,$\\mu$m observations of HD~107146 are consistent with previous spatially resolved disk observations at 1.3\\,mm \\citep{cor09} and in the near-infrared \\citep{ard04}, although we demonstrate that the long-wavelength data are distributed smoothly to within the precision of the observations.  For both systems, there is some indication that the population of dust grains giving rise to the millimeter data are located closer to the star than the likely smaller populations probed by the short-wavelength data, consistent with theoretical predictions for collisional destruction leading to blow-out of small dust.  Long-wavelength observations with higher sensitivity will be particularly exciting for the case of HR~8799, since they will be able to constrain the disk geometry (inclination and position angle) and more precisely determine the location of the inner edge of the dust belt, both of which will place important dynamical constraints on the masses of the directly-imaged planets in the system."
    },
    "1107/1107.4897_arXiv.txt": {
        "abstract": "The maximally extended Reissner--Nordstr\\\"{o}m (RN) manifold with $e^2 < m^2$ begs for attaching a material source to it that would preserve the infinite chain of asymptotically flat regions and evolve through the wormhole between the RN singularities. So far, the attempts were discouraging. Here we try one more possible source -- a solution found by Ruban in 1972 that is a charged generalisation of an inhomogeneous Kantowski--Sachs-type dust solution. It can be matched to the RN solution, and the matching surface must stay all the time between the two RN event horizons. However, shell crossings do not allow even half a cycle of oscillation between the maximal and the minimal size. ",
        "introduction": "Avoiding the Big Bang singularity in cosmological models had been a recurring idea in the literature. The hopes for constructing a model without singularity were largely dashed by the singularity theorems of Hawking and Penrose (see Ref. \\cite{HaEl1973} for a review). They were temporarily revived by the finding of Vickers \\cite{Vick1973} that in a charged dust model generalising that of Lema\\^{\\i}tre \\cite{Lema1933} and Tolman \\cite{Tolm1934} (LT) the Big Bang can be prevented by the charge distribution,\\footnote{This charged generalisation of the LT model was first found as a solution of the Einstein -- Maxwell equations by Markov and Frolov in 1970 \\cite{MaFr1970}.} provided that the absolute value of the charge density is, in geometrical units, \\textit{smaller} than the mass density. This finding created another expectation, that a finite ball of charged dust matched to the Reissner \\cite{Reis1916} -- Nordstr\\\"{o}m \\cite{Nord1918} (RN) solution would be able to collapse and bounce through the wormhole of the maximally extended RN manifold, thereby giving physical meaning to the infinite chain of black holes and asymptotically flat regions.\\footnote{The extension of the RN solution composed of two coordinate patches was first calculated by Graves and Brill in 1960 \\cite{GrBr1960}. The infinite mosaic of conformal diagrams shown in Fig. \\ref{reinormax} first appeared in the paper by Carter \\cite{Cart1966} in 1966, and was described in detail in a review article by Carter \\cite{Cart1973}. See Ref. \\cite{PlKr2006} for another pedagogical introduction.} However, this renewed hope was dashed again in 1991 by Ori \\cite{Ori1991} who proved that, under exactly the same conditions that prevent the Big Bang, shell crossings will inevitably appear and destroy the dust ball before it enters the wormhole. Then, Krasi\\'nski and Bolejko \\cite{KrBo2006,KrBo2007} found a gap in Ori's assumptions and tried to improve upon his result. Namely, Ori assumed that the ratio of the absolute value of the charge density to the mass density (call this ratio $\\alpha$) is smaller than 1 everywhere, including the centre of symmetry. Krasi\\'nski and Bolejko considered the case when $\\alpha \\to 1$ at the centre, while being $< 1$ everywhere else. The situation turned out to be better, but not definitively. With initial conditions carefully tuned, the charged dust ball could go through the wormhole just once, being destroyed by shell crossings soon after reaching the maximal size in the next asymptotically flat region. Moreover, a direction-dependent singularity necessarily appeared at the centre of the ball at the instant of minimal size. This is not a satisfactory situation, but the best result achieved so far.  In an attempt to find a better model, we now tried to match the Ruban \\cite{Ruba1972} charged dust solution to the RN metric and see what results. The charged Ruban solution is a generalisation to nonzero charge of the nonstatic dust solution investigated earlier by Ruban \\cite{Ruba1968,Ruba1969}, which in turn is an inhomogeneous generalisation of the Kantowski -- Sachs model \\cite{KaSa1966}.\\footnote{The neutral dust solution investigated by Ruban was first found as a solution of the Einstein equations by Datt in 1938 \\cite{Datt1938}, but instantly dismissed as being of ``little physical significance''.} The matching of these two solutions is very easily achieved, with an interesting result. The Ruban solution represents a pulsating charged dust ball whose outer surface must remain forever between the two event horizons of the RN solution. It touches the inner event horizon at its minimal size and the outer horizon at its maximal size. However, a simple investigation shows that shell crossings are inevitable inside the ball and do not allow even a half cycle of an oscillation, appearing between the maximum and minimum size both in the expansion phase and in the collapse phase.  Thus, the field is still open for finding a physical nonstatic source for the RN solution that would proceed through the wormhole. The next thing to do is to find a charged perfect fluid solution of the Einstein -- Maxwell equations in which pressure gradients would prevent the formation of shell crossings. This, however, promises to be an extremely difficult task.  In this paper, we first present the RN solution in coordinates adapted to the matching, then the charged Ruban solution as found by the author, then we prove that the matching can be done, and finally we prove that shell crossings in the Ruban solution are inevitable. We also investigate the limit of zero charge of the Ruban -- RN configuration. Then the dust source is matched to the Schwarzschild solution across a hypersurface that remains inside the Schwarzschild horizon, only touching it from inside at the moment of maximal expansion. Shell crossings can then be avoided by an appropriate choice of the arbitrary functions in the Ruban model. However, this model has a finite time of existence, being born in a Big Bang that matches to the past Kruskal \\cite{Krus1960} -- Szekeres \\cite{Szek1960} singularity and crushing into a Big Crunch that matches to the future KS singularity.  We use the signature $(+ - - -)$. We do not assume ``units in which $G = c =1$''; whenever these constants are absent, this means they were absorbed into other symbols. For example, our time coordinate will be $t = c \\times$ [the physical time], our $M$ will denote $(G / c^2) \\times$ [the mass], etc. The labelling of the coordinates is $(x^0, x^1, x^2, x^3) = (t, r, \\vartheta, \\varphi)$. ",
        "conclusions": "\\label{summ}  \\setcounter{equation}{0}  We investigated the charged Ruban solution \\cite{Ruba1972} as a possible matter source for the maximally extended Reissner -- Nordstr\\\"{o}m solution with $e^2 < m^2$. The matching of these two solutions is easily achieved. The hypersurface that forms the boundary between the Ruban and RN regions stays all the time between the two RN event horizons, touching the inner one at its minimal radius and the outer one at maximal radius. However, shell crossings make the completion of even half a cycle of such an oscillation impossible; they appear between each $[r_-, r_+]$ and each $[r_+, r_-]$ pair.  Shell crossings can be avoided in the limit of zero charge, when the arbitrary functions in the Ruban solution obey a simple inequality. Then the Ruban solution can be matched to the Schwarzschild solution inside the event horizon, but the model has a finite time of existence.  The problem of finding a matter source to the maximally extended RN solution thus still remains open, with one more possibility being now elliminated.  \\appendix"
    },
    "1107/1107.4290_arXiv.txt": {
        "abstract": "We model the evolution of infrared galaxies using a phenomenological approach to match the observed source counts at different infrared wavelengths. In order to do that, we introduce a new algorithm for reproducing source counts which is based on direct integration of probability distributions rather than using Monte-Carlo sampling. We also construct a simple model for the evolution of the luminosity function and the colour distribution of infrared galaxies which utilizes a minimum number of free parameters; we analyze how each of these parameters is constrained by observational data. The model is based on pure luminosity evolution, adopts the Dale \\& Helou Spectral Energy Distribution (SED) templates, allowing for evolution in the infrared luminosity-colour relation, but does not consider Active Galactic Nuclei as a separately evolving population. We find that the $850\\,\\mu$m source counts and redshift distribution depend strongly on the shape of the luminosity evolution function, but only weakly on the details of the SEDs. Based on this observation, we derive the best-fit evolutionary model using the $850\\,\\mu$m counts and redshift distribution as constraints.  Because of the strong negative $K$-correction combined with the strong luminosity evolution, the fit has considerable sensitivity to the sub-$L_*$ population at high redshift, and our best-fit shows a flattening of the faint end of the luminosity function towards high redshifts. Furthermore, our best-fit model requires a colour evolution which implies the typical dust temperatures of objects with the same luminosities to decrease with redshift. We then compare our best-fit model to observed source counts at shorter and longer wavelengths which indicates our model reproduces the $70\\,\\mu$m and $1100\\,\\mu$m source counts remarkably well, but under-produces the counts at intermediate wavelengths.  Analysis reveals that the discrepancy arises at low redshifts, indicating that revision of the adopted SED library towards lower dust temperatures (at a fixed infrared luminosity) is required. This modification is equivalent to a population of cold galaxies existing at low redshifts, as also indicated by recent Herschel results, which are underrepresented in IRAS sample. We show that the modified model successfully reproduces the source counts in a wide range of IR and submm wavelengths. ",
        "introduction": "Although dust is an unimportant component in the mass budget of galaxies, its presence radically alters the emergent spectrum of star forming galaxies. Since stars are born in dusty clouds, most of the energy of young stars is absorbed by dust particles, which are heated by the absorption process and radiate their energy away by thermal emission at infrared (IR) and submillimetre (submm) wavelengths. As a result, the infrared radiation of star forming galaxies is a useful measure of the massive star formation rate. Mainly due to the atmospheric opacity, the thermal radiation from dust could not be systematically studied until the launch of IRAS in the mid-1980s, which provided the first comprehensive view of the nearby dusty Universe.  Ten years later COBE discovered the Cosmic Infrared Background (CIB) \\citep{Puget96,Fixsen98} and it turned out the observed power of the CIB is comparable to what can be deduced from the optical Universe. This was in contrast to the observations of local galaxies which suggest only one third of the energy output of galaxies is in IR bands \\citep{Lagache05}.  Moreover, the relatively flat slope of the CIB at long wavelengths indicated a population of dusty galaxies which are distributed over a wide range of redshifts \\citep{Gispert00}.  Thanks to various large surveys performed with different satellites and ground based observatories, this background radiation is now partly resolved into point sources at different wavelengths. While at shorter wavelengths the emission is mainly coming from local and low redshift galaxies, longer wavelengths contain information about larger distances.  The shape of emission spectrum of warm dust particles resembles a modified blackbody spectrum with a peak varying with the typical dust temperature which is observed to be around $T\\sim 30-40$K in galaxies; therefore the far-IR (FIR) and submm spectrum of a typical dusty galaxy consists of a peak at rest-frame wavelengths around $\\lambda \\sim 100-200 \\mu m$ which drops on both sides (see Figure~\\ref{fig:SED}). While the presence of other emitters like polycyclic aromatic hydrocarbons (PAHS) and Active Galactic Nuclei (AGNs) complicates the shape of spectra at shorter wavelengths, at submm wavelengths the spectra behave like modified blackbodies and their amplitudes drop steeply. In fact, because of this steep falloff in the Rayleigh-Jeans tail of the Spectral Energy Distribution (SED), the observed flux density at a fixed submm wavelength can rise by moving the SED in redshift space (the so called $K$-correction), which counteracts the cosmological dimming. It turns out that the interplay between these two processes enables us to observe galaxies at submm wavelengths out to very high redshifts (see also the discussion in Section~\\ref{sec:others} and Figure~\\ref{fig:short-long-fluxes}).   After the first observations of SCUBA at $850\\,\\mu$m confirmed the importance of submm galaxies at high redshifts \\citep{Smail97}, there were many subsequent surveys using different instruments which explored different cosmological fields \\citep{Smail02,Webb03,Borys03,Greve04,Coppin06,Bertoldi07,Austermann10,Scott10,Vieira10} and also some surveys which used gravitational lensing techniques to extend the observable submm Universe to sub-mJy fluxes \\citep{Smail97,Chapman02,Knudsen08, Johansson11}.  While low angular resolution makes individual identifications and spectroscopy a daunting task, the surface density of sources as a function of brightness (i.e., the source counts) can be readily analysed and contains significant information about the population properties and their evolution. One can assume simple smooth parametrized models for the evolution of dusty galaxies and relate their low redshift observed properties (e.g., the total IR luminosity function, which we simply call the luminosity function, or in short LF hereafter) to their source counts \\citep{Blain93,Guiderdoni97,Blain99,Chary01,Rowan-Robinson01,Dole03,Lagache04,Lewis05,LeBorgne09,Valiante09}. Such backward evolution models usually combine SED templates and the low redshift properties of IR galaxies which are allowed to change with redshift based on an assumed parametric form, together with Monte-Carlo techniques to reproduce the observed source counts. These models are therefore purely phenomenological, and suppress the underlying physics. Their power lies exclusively in providing a parametrized description of statistical properties of the galaxy population under study, and the evolution of these properties. The results of such modeling provides a description of the constraints that must be satisfied (in a statistical sense) and could be explored further by more physically motivated models such as hydrodynamical simulations embedded in a $\\Lambda$CDM cosmology, with subgrid prescriptions for star formation and feedback (e.g. \\citet{Schaye10}). It is important to realize that backward evolution models are constrained only over limited observational parameter space (e.g., at particular observing wavelengths, flux levels, redshift intervals, etc.), and are ignorant of the laws of physics. Therefore it is dangerous to use them outside of their established validity ranges; in other words, these models have descriptive power but little explanatory power. As a result, one of the most instructive uses of these models is in analyzing where they fail to reproduce the observations, and how they can be modified to correct for these failures which implies that the underlying assumptions must be revised. This is the approach that we use in the present paper.  Since nowadays many and various observational constraints exist (e.g. source counts at various IR and submm wavelengths, colour distributions, redshift distribution), backward evolution models can reach considerable levels of sophistication.  However, early backward evolution models used only a few SED templates (or sometimes even only one SED) to represent the whole population of dusty galaxies. This approach neglects the fact that dusty galaxies are not identical and span a variety of dust temperatures and hence SED shapes. As a result, such models cannot probe the possible evolution in the SED properties of submm galaxies, for which increasing observational support has been found during the last few years \\citep{Chapman05,Pope06,Chapin09,Symeonidis09,Seymour10,Hwang10}.  In addition, existing backward evolution models typically use only ``luminosity'' and/or ``density'' evolution which respectively evolves the characteristic luminosity (i.e., $L_\\star$) and the amplitude of the LF without changing its shape. In other words, they do not consider any evolution in the shape of the luminosity function. Finally, the intrinsically slow Monte-Carlo methods used in many of the models make the search of parameter space for the best-fit model laborious and tedious.  In this paper, we use a new and fast algorithm which is different from conventional Monte-Carlo based algorithms, for calculating the source counts for a given backward evolution model.  We also use a complete library of IR SED templates which form a representative sample of observed galaxies, at least at low redshifts. This will enable us to address questions about the evolution of colour distribution of dusty galaxies during the history of the Universe. We also consider a new evolution form for the LF which allows us to constrain evolution in the shape of the LF in addition to quantifying its amplitude changes.  The structure of the paper is as follows: in \\S\\ref{sec:ingredients} we present different ingredients of our parametric colour-luminosity function (CLF) evolution model and introduce a new algorithm for the fast calculation of the source counts for a given CLF. Then, after choosing our observational constraints for $850\\,\\mu$m objects in \\S\\ref{sec:observed} we try to find a model consistent with those constraints by studying the sensitivity of the produced source counts to different parameters in \\S\\ref{sec:modelprop}. After finding a $850\\,\\mu$m-constrained best-fit model, we test its performance in producing source counts at other wavelengths in \\S\\ref{sec:others} which leads us to a model capable of producing source counts in a wide range of IR and submm wavelengths. Finally, after discussing the astrophysical implications of our best-fit model in \\S\\ref{sec:discussion}, we end the paper with concluding remarks. Throughout this paper, we use the standard $\\Lambda$CDM cosmology with the parameters $\\Omega_{\\rm{m}}=0.3$, $\\Omega_{\\rm{\\Lambda}}=0.7$ and $h=0.75$. ",
        "conclusions": "\\label{conclusion}  We have described a backward evolution model for the IR galaxy population, with a small number of free parameters, emphasizing which parameters are constrained by which observations.  We also introduced a new algorithm for calculating source counts for a given evolutionary model by direct integration of probability distributions which is faster than using Monte-Carlo sampling. This is an important advantage for searching large volumes of parameter space for the best-fit model.  While most of the earlier works used only one or a handful of SED templates to represent the whole population of IR objects, we used a library of IR SEDs which are able to match the IR properties of the large variety of observed star-forming objects. This approach is necessary in order to model the colour evolution of IR galaxies in addition to produce simultaneously the counts and the redshift distributions at wavelengths shorter than $850\\,\\mu$m.  Contrary to some other models, we assumed a negligible contribution from AGN in our SED templates, noting the inclusion of AGN is only necessary for reproducing the properties of IR galaxies at very short IR wavelengths\\footnote{For instance at $24\\,\\mu$m where our model under-produces the counts at flux thresholds $0.1<S_{\\rm{24}} < 10{\\rm{mJy}}$ by a factor of $\\sim 1.5-2$ but matches the observed data at fainter and brighter fluxes} which could also be sensitive to other modeling difficulties such as the PAH contribution to the SEDs.  We used available $850\\,\\mu$m source counts together with the redshift distribution of submm galaxies to constrain our best-fit model. At $850\\,\\mu$m, due to the K-correction, the source count is sensitive to the evolution of IR galaxies in a wide redshift range and out to very high redshifts. We showed that there is a degeneracy between the rate by which the characteristic luminosity of IR galaxies should increase to reproduce the source count and the maximum redshift out to which this increase should be continued; we resolve this degeneracy by requiring the model to reproduce the observed redshift distribution of submm galaxies. Moreover, we showed that our model requires a colour evolution towards cooler typical dust temperatures at higher redshifts. The employed colour evolution is similar to that used by \\citet{Valiante09} however, our best-fit model predicts a somewhat stronger colour evolution than that proposed by these authors.  Another important feature of our model is that the best-fit is obtained using pure luminosity evolution but mildly evolving high-luminosity and low-luminosity slopes in the LF\\null. Since high-luminosity sources are rare, the evolution in the high-luminosity slope is of little consequence. However, the evolution of the low-lumuninosity slope affects large numbers of galaxies and if confirmed, this effect must have a physical origin, which can be addressed using numerical simulations of the evolution of the galaxy population, as well as with a combination of deep Herschel and optical imaging.  The $850\\,\\mu$m-constrained best-fit model is consistent with observed $1100\\,\\mu$m and $70\\,\\mu$m source counts, which confirm the consistency of the implied colour-lumionosity-redshift distribution at both low and high redshifts. However, this model under-produces the observed source counts at intermediate wavelengths, namely at $160\\,\\mu$m and SPIRE bands. To resolve this issue, we used the observed data at different wavelengths to find best-fit models which can reproduce their observed source counts. While the best-fit models constrained at $70\\,\\mu$m and $850\\,\\mu$m are consistent with each other and also $1100\\,\\mu$m, the implied evolutions for models capable of reproducing observed counts at other wavelengths are too diverse to be reconciled in a single model; specifically they need too strong colour evolutions which contradicts $850\\,\\mu$m observations. While the inconsistency at $160\\,\\mu$m has been reported in earlier works \\citep{LeBorgne09,Valiante09}, we are the first to report it for $250\\,\\mu$m, $350\\,\\mu$m and $500\\,\\mu$m. We showed that the source counts at these wavelengths can be reproduced consistently, by adopting the best-fit model which produces correct $70\\,\\mu$m, $850\\,\\mu$m and $1100\\,\\mu$m source counts, together with a modification in SED templates which is equivalent to the existence of a cold population of dusty galaxies at low to intermediate redshifts which are under-represented in IRAS data. Besides the fact that there is some observational evidence for the existence of such galaxies \\citep{Stickel98,Stickel00,Chapman02,Patris03,Dennefeld05,Sajina06,Amblard10}, this assumption is similar to what other models had to assume in order to reproduce adequate $160\\,\\mu$m sources \\citep{Lagache03,Lagache04,Valiante09}.  It is important to keep in mind that phenomenological models like what we described in this paper, are mainly simple mathematical forms which relate different observations consistently rather than being physical models with explanatory power. However, their performance at different wavelengths and the distribution of sources they require for different redshifts can be used as their main predictions which also could be used to test their validity. While we used the redshift distribution of submm galaxies to constrain our model, we note that the observed redshift distribution of other wavelengths, if available, are in agreement with our best-fit model predictions \\citep{Jacobs11,Berta11}.  Additional information including tabulated data for differential and cumulative source counts at different wavelengths and their redshift distributions is available at $\\mathtt{http://www.strw.leidenuniv.nl/genesis/}$"
    },
    "1107/1107.0002.txt": {
        "abstract": "We make use of interferometric CO and \\hi\\ observations, and optical integral-field spectroscopy from the \\atlas\\ survey to probe the origin of the molecular and ionised interstellar medium (ISM) in local early-type galaxies. We find that 36$\\pm$5\\% of our sample of fast rotating early-type galaxies have their ionised gas kinematically misaligned with respect to the stars, setting a strong lower limit on the importance of externally acquired gas (e.g. from mergers and cold accretion). Slow rotators have a flat distribution of misalignments, indicating that the dominant source of gas is external. The molecular, ionised and atomic gas in all the detected galaxies are always kinematically aligned, even when they are misaligned from the stars, suggesting that all these three phases of the interstellar medium share a common origin. In addition, we find that the origin of the cold and warm gas in fast-rotating early-type galaxies is strongly affected by environment, despite the molecular gas detection rate and mass fractions being fairly independent of group/cluster membership. Galaxies in dense groups and the Virgo cluster  nearly always have their molecular gas kinematically aligned with the stellar kinematics, consistent with a purely internal origin (presumably stellar mass loss). In the field, however, kinematic misalignments between the stellar and gaseous components indicate that at least 42$\\pm$5\\% of local fast-rotating early-type galaxies have their gas supplied from external sources. When one also considers evidence of accretion present in the galaxies' atomic gas distributions, $\\gtsimeq$46\\% of fast-rotating field ETGs are likely to have acquired a detectable amount of ISM from accretion and mergers. We discuss several scenarios which could explain the environmental dichotomy, including preprocessing in galaxy groups/cluster outskirts and the morphological transformation of spiral galaxies, but we find it difficult to simultaneously explain the kinematic misalignment difference and the constant detection rate. Furthermore, our results suggest that galaxy mass may be an important independent factor associated with the origin of the gas, with the most massive fast-rotating galaxies in our sample (M$_K\\ltsimeq-24$ mag; stellar mass of $\\approx$8$\\times$10$^{10}$ M$_{\\odot}$) always having kinematically aligned gas. This mass dependence appears to be independent of environment, suggesting it is caused by a separate physical mechanism. ",
        "introduction": "Lenticular and elliptical galaxies, collectively known as early-type galaxies (ETGs), have long been thought to be the endpoint of galaxy evolution. These systems have surprisingly uniform red optical colours and are located in a tight red sequence in an optical colour-magnitude diagram \\citep[e.g.][]{Baldry:2004p3398}. The origin of the colour bimodality that separates the ETGs in the red sequence from star-forming galaxies, located in a `blue cloud', has become a central question in extragalactic astrophysics.  In order for galaxies to fade and join this tight red sequence rapidly enough, it is thought that the fuel for star formation must be consumed, destroyed or removed on a reasonably short timescale \\citep[e.g.][]{Faber:2007p3453}.  In the paradigm of cold dark matter with a non-zero cosmological constant (\\lcdm), it is thought that star-formation is shut-off in these galaxies by major mergers that can funnel the cold gas into the centre of the remnant \\citep{Barnes:2002p2041}, where it may be destroyed by an active galactic nucleus (AGN) or be used up quickly in a starburst. In massive galaxies, virial shocks \\citep[such as those discussed by][]{Birnboim:2003p3427} can also create a halo mass threshold, above which incoming gas will always shock to the virial temperature, stopping the accretion of further cold gas. Energy input from an AGN or starburst may also remove cold gas without a major merger, and keep halo gas from cooling. Galaxies in clusters and groups may also have their interstellar medium (ISM) removed by harassment, or stripped and ablated away by a hot intra-cluster medium (ICM).  Evolutionary processes like those listed above leave galaxies on the red sequence with little or no cold ISM, and thus no star formation. Evidence is mounting, however, that a reasonable fraction of ETGs do have cold gas reservoirs and residual star-formation. For example the surveys of \\cite{Grossi:2009p3448}, \\cite{Morganti:2006p1934} and \\cite{Oosterloo:2010p3376} find between 44 and 66\\% of field ETGs have detectable neutral hydrogen reservoirs, while \\cite{Colbert:2001p1791} have shown that over 75\\% of early-types contain dust features in optical images. \\cite{Yi:2005p3450} and \\cite{Kaviraj:2007p1804} have also shown that at least 30\\% of ETGs have an excess of UV emission, attributable to ongoing residual star formation (over and above that predicted for an old stellar population from the UV-upturn phenomenon; see \\citealt{Yi:2008p3435} for a review). \\cite{Sarzi:2006p1474} have shown that at least 10\\% of local ETGs have strong Balmer line emission and line ratios indicative of ongoing star-formation. These studies all imply some ETGs must have a cold molecular gas reservoir to fuel this ongoing star-formation.  The molecular gas content of early-types had until recently been investigated mainly using infrared-selected surveys. For example \\cite{Knapp:1996p1859} report a molecular gas detection rate of 80\\% for ETGs brighter than 1 Jy at 100 $\\mu$m. \\cite{Combes:2007p231} performed one of the first studies of ETGs which was not infrared-selected. They observed the SAURON sample of galaxies {(a representative sample of 48 ETGs which had been observed with the SAURON integral field spectrograph, see \\citealt{deZeeuw:2002p1496} for full details)} and reported a molecular gas detection rate of 28\\% (with a CO(1-0) sensitivity of 3mK T$_a^{*}$ in a 30 \\kms\\ channel). {The survey of \\cite{Welch:2010p3290} \\citep[which combined and added to the results of][]{Welch:2003p2521,Sage:2006p3466,Sage:2007p3467}} was the first published volume-limited sample, which found a detection rate of 26\\% for ETGs within 20 Mpc (with a similar sensitivity).  The \\atlas\\ project is a complete, volume-limited survey of the properties of 260 morphologically-selected ETGs within 40 Mpc \\citep[for full details see][hereafter Paper I]{Cap2010}. This survey combines integral-field spectroscopy (IFS), optical imaging, molecular and atomic gas observations and simulations. As part of this survey \\cite{Young2010} (hereafter Paper IV) have presented an unbiased census of the molecular gas content of this sample of nearby early-types, and report that 22\\% of optically-selected, morphologically-classified ETGs have substantial molecular gas reservoirs (10$^7$--10$^9$ M$_{\\odot}$ of H$_2$). The brightest detections have been followed up with CO-interfeometry (Alatalo et al., in prep; see Section \\ref{coobs}). Serra et al., in prep report a \\hi\\ detection rate of $\\approx$40\\% (for field galaxies with declination above 10$^{\\circ}$ drawn from the same sample). These studies are the best available complete, unbiased surveys of the ISM in local ETGs, showing that a reasonable fraction of these systems have sizable gas reservoirs.  If galaxies form in a hierarchical manner, then these observations pose a challenge to the standard view that early-type galaxies join and remain on the red-sequence due to a lack of cold gas. One must either demonstrate that it is possible to create a tight red sequence without removing all of the cold ISM, or that galaxies can regenerate or acquire cold gas after joining the red sequence.  The gas we detect in these galaxies could be a remnant, left over from the progenitors that formed the ETG. Stellar population analyses (e.g. \\citealt{Kuntschner:2006p1489}, \\citealt{Kuntschner:2010p3471}, McDermid et al., in prep), however, show that the underlying stellar populations in these galaxies are old ($\\gg1$\\,Gyr). When combined with the fact that most ETGs show relaxed stellar kinematics \\citep[][hearafter Paper III]{Emsellem2010} and evidence from deep imaging (Duc et al., submitted; Paper IX) this suggests that the last major merger that could have formed such systems was at least several gigayears ago. Unless the remnant gas can be made stable against star-formation this thus suggests that we are seeing regenerated or newly accreted gas.  Galaxies can regenerate a cold ISM through both \\textit{internal} and \\textit{external} processes. Stellar evolution models predict that an average stellar population returns on the order of half its stellar mass to the ISM over a Hubble time \\citep{Jungwiert:2001p3209,Lia:2002p3210,Pozzetti:2007p3211,Martig:2009p2862}. This \\textit{internal} gas return is dominated not by cataclysmic events such as supernovae, but by the long-term stellar mass loss from red giant branch, (post-)asymptotic giant branch stars and planetary nebulae \\citep{Parriott:2008p2869,Bregman:2009p2870}.  The majority of the ejected mass from giant stars will shock heat and join the hot gas reservoir of the galaxy \\citep{Parriott:2008p2869}.  It may then be possible for this material to cool, recombine, fall toward the galaxy centre and eventually become molecular, regenerating the cold gas reservoir. The fraction of gas which is able to cool from the hot halo in this way is however unknown. Mass loss from planetary nebulae may avoid being shock heated, and cool from the ionised phase directly without joining the galaxies hot halo \\citep{Bregman:2009p2870}. Simulations have shown that this mass loss looses some angular momentum in this process, but generally it preserves the general sense of rotation of the parent star \\citep[e.g.][]{Martig:2010p3475}.  Stellar mass loss must be occurring in all ETGs at all times, at a rate depending on the number of stars present in each galaxy and its star formation history. One would thus expect many ETGs to have detectable molecular gas (assuming that some fraction of the hot gas reservoir does cool). Paper IV however found that only 22\\% of ETGs have detectable molecular gas, and there is no dependence of the detection rate and molecular gas mass fraction on host luminosity. {It is of course possible that gas reservoirs exist in the other systems with a more quiescent recent star-formation history, resulting in low gas masses from stellar mass loss that are below our detection threshold.} The lack of correlation with the host galaxy properties, however, suggests that other mechanisms must be destroying or preventing the molecular gas from forming in 78\\% of ETGs, and that stellar mass loss may not be the dominant method of acquiring cold gas.  Galaxies can also regain a cold ISM via \\textit{external} processes, such as (major and minor) mergers and/or cold mode accretion from the intergalactic medium (IGM). In a minor merger or cold accretion some fraction of the gas can avoid being shocked and will fall to the centre of the galaxy and cool, often producing a central disc or rings \\citep{Mazzuca:2006p3379,ElicheMoral:2009p2918}. In a major merger ejected gas can be re-accreted over the course of cosmic time. \\hi\\ in ETGs has long been thought to come almost exclusively from external sources \\citep[e.g][]{Knapp:1985p3440}. More recently deep HI observations \\citep[e.g.][]{Morganti:2006p1934} have confirmed that almost all \\hi-detected field ETGs show signs of accretion, indicating an external origin for much of the atomic gas, with an average \\hi\\ accretion rate of 0.1 M$_{\\odot}$ yr$^{-1}$ \\citep{Oosterloo:2010p3376}. If the same is true for the molecular and ionised gas has yet to be determined.   Recent work has suggested one way to completely sidestep this problem, via `morphological quenching'  \\citep{Martig:2009p2923}. In this scenario, red sequence galaxies are created not by removing all the cold ISM, but by making some fraction of it stable against gravitational collapse. The growth of a stellar spheroid via merging or secular processes increases the stellar density, and hence the epicyclic frequency and velocity dispersion in the centre of the galaxy. This can cause the Toomre stability criterion \\citep{Toomre:1964p3272} to exceed unity, implying that the gas is stable against gravitational collapse. As the star formation rate would be very low, this stable cold ISM would persist for many billions of years, and hence would be detectable today in red galaxies. In some \\atlas\\ galaxies we do indeed see large amounts of \\hi\\ in a regular disk, but no corresponding young stellar population \\citep{Morganti:2006p1934}. This process cannot be {how every ETG transits to the red sequence}, however, as such long-lived discs are not observed in all ETGs.  In summary, determining the dominant source of the cold ISM in ETGs is vital in order to understand their formation and evolution. If stellar mass loss can build up molecular reservoirs, then galaxies can transform themselves from spheroidal to disky systems over time and in isolation, and perhaps even (if morphological quenching does not intervene) evolve back into the blue cloud over cosmic timescales \\citep[e.g.][]{Kannappan:2009p3436}. If however mergers are the dominant source of the gas, then star-formation episodes are likely to be short-lived and possibly violent, and the gas will not accumulate and regenerate.  Observationally the origin of the  gas may be addressed by comparing the angular momentum of the ISM with that of the underlying stellar population. Because of angular momentum conservation, stellar mass loss must produce gas that is kinematically aligned with the bulk of the stars which produced it. To first order, material from external sources can enter a galaxy with any angular momentum, so many such mismatches should be observable today, the exact fraction depending on the dynamical timescales for misaligned material to relax into the equatorial plane. Even in very old systems, one might thus expect to see an equal number of galaxies with counter and co-rotating gas if external accretion and mergers are important (and if there is no preferred accretion direction).  A number of previous studies have used the misalignment of the angular momenta of gas and stars to constrain the origin of the ionized gas \\citep[e.g.][]{Kannappan:2001p3451,Sarzi:2006p1474}, and the molecular gas \\citep[e.g.][]{Young:2002p943,Young:2008p788,Crocker:2011p3291} in early-type galaxies. The goal of this paper is to use the \\atlas\\ complete volume limited sample of 260 ETGs to constrain the importance of externally acquired gas in this way.    In Section 2 of this paper we describe the observations our work is based upon, and present our method for extracting kinematic position angles. In Section 3 we describe the misalignment distributions obtained, and we discuss them further in Section 4. We conclude in Section 5 and discuss prospects for the future. ",
        "conclusions": "In this work we have demonstrated that a large proportion of the gas found in local early-type galaxies is likely supplied by external processes. Gas kinematically misaligned with respect to the stars is common, indicating an external origin for the gas in $\\gtsimeq$36\\% of all fast-rotating ETGs. We have also shown that the ionised, atomic and molecular gas in local ETGs are linked, always having similar kinematics and thus presumably sharing a common origin. In the field, $\\approx$42\\% of galaxies have kinematically misaligned ionised gas, confirming that mergers and accretion play a vital role in supplying gas to field ETGs. Slow rotators are not generally detected in molecular gas, but show a flat distribution of ionised gas kinematic misalignments, suggesting they obtain ionised gas primarily through mergers and/or accretion. In many cases, fast-rotating galaxies with kinematically aligned molecular and ionised gas have \\hi\\ distributions that again suggest an external origin for the gas. This suggests that external gas sources dominate in the field (as has been suggested by previous authors).   In the Virgo cluster, the molecular and ionised material in fast-rotators is nearly always kinematically aligned with the bulk of the stars, pointing to gas supplied by purely internal processes. Fast-rotators in dense groups also appear to always have aligned gas kinematics, suggesting that the local environment is important in understanding the environmental dependance of the gas origins. Given the results of Paper IV, indicating similar molecular gas detection rates and mass fractions in clusters and the field, this sets constraints on potential mechanisms for generating and influencing molecular and ionised gas in cluster environments.  Molecular and ionised gas in the most massive fast-rotating galaxies also appear to always be kinematically aligned with the stars, independent of environment, suggesting that the alignment can be caused by galaxy scale processes which reduce the probability that cold, kinematically misaligned gas can be accreted onto the galaxy (e.g. AGN feedback, the ability to host a hot X-ray gas halo, or a halo mass threshold). This also supports the picture (discussed in Papers III and VIII) that fast-rotating galaxies (which are always aligned at high mass) have a different formation mechanism to slow-rotators (who at similarly high masses mainly have externally supplied ionised gas).     We tentatively suggest that preprocessing in groups may help to explain the environmental dichotomy in the origin of the gas. Merger-induced starbursts in groups and cluster outskirts could consume any kinematically misaligned molecular gas that is present, and leftover ionised material would be ram pressure-stripped. Once galaxies settle into a cluster or \\hi-poor group, external gas accretion and mergers are suppressed, allowing stellar mass loss to regenerate a kinematically aligned gas reservoir.  It is possible that features such as bars and rings could funnel dust and gas lost by stars to the centre of the galaxy, or collect them together, and could explain the greater efficiency with which galaxies in dense environments must recreate their dense gas reservoirs. Alternatively, we could be detecting the remnant gas left over from the morphological transformation of spiral galaxies into ETGs as they enter the cluster and group environments. Both of these possibilities, however, fail to explain why the detection rate of molecular gas (and the molecular gas mass fractions) are similar inside and outside of the Virgo cluster.  More work to pin down the mechanism(s) creating and/or aligning the gas is clearly required to unambiguously determine which, if any, of these effects is dominant.  It would also be beneficial to extend this sort of analysis to other clusters and groups, to see if the results reported hold true in yet denser environments. The Fornax cluster in the southern hemisphere and the Coma cluster in the north are obvious nearby targets, which should become accessible once new facilities such as the Large Millimeter Telescope (LMT) and the Atacama Large Millimeter/sub-millimeter Array (ALMA) come online.    \\vspace{10pt} \\noindent \\textbf"
    },
    "1107/1107.2013_arXiv.txt": {
        "abstract": "The standard axion-like particle explanation of the observed large-scale coherent orientations of quasar polarisation vectors is ruled out by the recent measurements of vanishing of circular polarisation. We introduce a more general wave-packet formalism and show that, although decoherence effects between waves of different frequencies can reduce significantly the amount of circular polarisation, the axion-like particle hypothesis is disfavoured given the bandwidth with which part of the observations were performed. Finally, we show that a more sophisticated model of extragalactic fields does not lead to an alignment of polarisations. ",
        "introduction": "A frequent prediction of extensions of the Standard Model of particle physics is the existence of stable weakly interacting light (sub-eV) scalar or pseudoscalar particles. The `invisible' axion~\\cite{Kim:1979,SVZ:1979,Zhitnitsky:1980,DFS:1981,Weinberg:1978,Wilczek:1978} is certainly the best-known candidate, so that any particle of this kind is nowadays commonly referred to as an {\\emph{axion-like particle}} (ALP) ---even though it might have nothing to do with the Peccei--Quinn solution to the strong CP problem\\cite{PecceiQuinn:1977}. Usually, the smallness of the masses of these particles is related to a very-high-energy scale where there would be new physics. Among these ALPs, one finds for instance chameleons, coming from $f(\\!\\!\\!~R)$ theories, but also scalar and pseudoscalar particles from Kaluza-Klein theories, super strings, or other theories beyond the Standard Model, which could be testable predictions; for recent reviews, see for instance~\\cite{Raffelt:2006rj,Jaeckel:2010ni} and references therein.  The ALPs can have a coupling to photons, as in the case for the axion~\\cite{Sikivie:1983ip,Raffelt:1987im}. As the information we get from astrophysics comes mainly from photons from distant sources, this property makes them an appealing ingredient in many astrophysical models, and their existence could be probed by astrophysical observations.  In fact, several authors have already reported different phenomena that might find a common explanation if one supposes the existence of nearly massless axion-like particles (of mass $m\\le 10^{-10}$~eV, and of coupling to photons $g\\sim 10^{-11}$~GeV$^{-1}$), \\textit{e.g.} the transparency of the Universe to high-energy photons~\\cite{DeAngelis:2007dy}, the luminosity relations for active galactic nuclei (AGN) at different wavelengths~\\cite{Burrage:2009mj}, or the high-energy cosmic rays from blazars~\\cite{Fairbairn:2009zi}. \\bigskip  Another interesting observation has to do with the distribution of position angles for polarisation of visible light coming from quasars\\footnote{Hereafter, `quasar' stands for `high-luminosity AGN'.}. These angles indicate the direction of maximum polarisation for each source with respect to an arbitrary direction, usually the north equatorial pole. It has been reported that the distribution of these individual preferred directions in extremely large regions of the sky ($\\sim$~Gpc) is not random\\cite{Hutsemekers:1998, Hutsemekers:2001,Jain:2003sg, Hutsemekers:2005}. From the latest sample available (355 quasars), global statistical tests indicate that the probability for the observed distribution to be random is between $3~10^{-5}$ and $2~10^{-3}$, depending on the test applied\\cite{Hutsemekers:2005}. This observation is remarkable as there is {\\it a priori} no reason why one should expect such correlations over cosmological distances, larger than the most extended structures presently known in the Universe.  This analysis also indicates that the effect is not likely to be explained by local causes (influence of our galaxy, dust, etc.) and suggests that it requires something more exotic. One might think that it comes from an alignment of quasar axes across the Universe. It is known~\\cite{Rusk:1985,Joshi:2007yf,Borguet:2007kb,Borguet:2008tn} that, for a given quasar, the direction of preferred polarisation is related to its morphology, so that a global alignment of the axes of quasars would lead to aligned polarisations. On the other hand, if one supposes that the objects themselves are aligned, the effect should be present in radio waves. However, a study\\cite{Joshi:2007yf} based on a sample of 4290 objects (52 of them being part of the sample~\\cite{Hutsemekers:2005}) has shown that there is \\emph{no} evidence for alignments in radio waves. It thus seems that this class of explanations is disfavoured.  On the other hand, it has been believed that these data could naturally be explained by the mixing of light with axion-like particles in background magnetic fields~\\cite{Jain:2002vx,Das:2004qka,Piotrovich:2008iy,Payez:2008pm,Hutsemekers:2008iv,Payez:2009kc,Payez:2009vi,Payez:2010xb,Agarwal:2009ic}: this would generate an alignment in visible light while leaving the polarisation of radio waves unaffected, as the mixing depends on energy. In the present paper, we show that this requires a very specific choice of magnetic fields, and that in general the alignment effect cannot be explained by the mixing. Furthermore, we shall see that, according to recent data~\\cite{Hutsemekers:2010fw}, the cause of this effect cannot be photon-ALP mixing, even for magnetic fields leading to an alignment.  In the next section, we introduce our notations and recall results for the polarisation of light described by plane waves due to the mixing of axion-like particles with photons. We then discuss why this cannot explain the data for the circular polarisation of quasars and present a wave-packet treatment of the mixing in Section~\\ref{sec:newcircpoldata}. Our analysis shows that, despite promising phenomenological implications, even wave-packets cannot reconcile the axion-like particle hypothesis with the full quasar sample. Further checks are made in Section~\\ref{sec:ruledout}, where we use different models for the magnetic field encountered by the incoming photons. ",
        "conclusions": "In this paper, we have considered photon-pseudoscalar mixing as the source of the observed alignments. This process has so far been the best hope to explain the observed data. The typical treatment assumes that the same faint coherent field is traversed by the light beams from all the quasars. Note that this is problematic, as the data display two alignments: quasars at redshift $z>1$ seem to be aligned in a different direction from closer ones. In order to obtain such an effect, one needs to assume a slice of coherent fields 1~Gpc large, with an intensity of the order of 0.1~$\\mu$G. Although this seems unlikely, one cannot rule out the explanation on those grounds, as so little is known about magnetic fields at high redshifts.  Recent data have shown that the circular polarisation of these objects is negligible, contrarily to the predictions from photon-pseudoscalar mixing. At face value, this kills the interpretation. However, we showed that it is possible to argue that, for white light, the average over frequencies suppresses the circular polarisation enough. From the sample of 355 quasars, only 6 circular polarisation measurements were obtained in white light, and 21 using a broadband Bessell V-filter, for which the bandwidth is nevertheless not broad enough to average the circular polarisation to almost zero.  Keeping the idea of averaging, we extended the model of magnetic field, by assuming a collection of cells of 100 kpc, with random 2~$\\mu$G fields averaging to 0.3~$\\mu$G, and letting most quantities fluctuate (size of cells, electron density, magnetic field). While these fluctuations can somewhat reduce the circular polarisation, they also destroy the alignment which was the first reason to consider photon-pseudoscalar mixing.  Hence it seems that the combination of the alignments and of the absence of circular polarisation remains at present a puzzle, which cannot be explained by photon-pseudoscalar mixing. One should note that this conclusion is based on the lack of circular polarisation in 19 objects measured with the V-filter. It is of course possible that magnetic field configurations are special for these 19 objects, and some more data~---or, equivalently, data in an even smaller bandwidth---~may be needed to reach certainty.  As a final note, the data on quasar polarisation show that the polarisations are only typically a few percent. Photon-pseudoscalar mixing can be much more efficient than this at producing polarisation. Hence the quasar data can be used to exclude part of the parameter space of ALPs. This will be the subject of a future paper~\\cite{tocome}."
    },
    "1107/1107.5813.txt": {
        "abstract": "{The unfortunate case where the Large Hadron Collider (LHC) fails to discover physics Beyond the Standard Model (BSM) is sometimes referred to as the ``Nightmare scenario'' of particle physics. We study the consequences of this hypothetical scenario for Dark Matter (DM), in the framework of the constrained Minimal Supersymmetric Standard Model (cMSSM). We evaluate the surviving regions of the cMSSM parameter space after null searches at the LHC, using several different LHC configurations, and study the consequences for DM searches with ton-scale direct detectors and the IceCube neutrino telescope. We demonstrate that ton-scale direct detection experiments will be able to conclusively probe the cMSSM parameter space that would survive null searches at the LHC with 100\\,fb$^{-1}$ of integrated luminosity at 14\\,TeV. We also demonstrate that IceCube (80 strings plus DeepCore) will be able to probe as much as $\\simeq17\\%$ of the currently favoured parameter space after 5\\,years of observation.} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{document}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "\\label{sec:intro} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  One of the main goals of the Large Hadron Collider (LHC) at CERN is to test the existence of physics Beyond the Standard Model (BSM), especially in view of the possible connection with the problem of Dark Matter (DM) in astrophysics and cosmology \\cite{Bertone:2010zz}. Recently, the ATLAS and CMS experimental collaborations have published the results of the first LHC searches for Supersymmetry (SUSY) \\cite{Chatrchyan:2011wc,Khachatryan:2011tk,daCosta:2011qk,Aad:2011hh, Aad:2011ks,Aad:2011xm}, based on  $\\sqrt{s} = 7$\\,TeV collisions recorded during 2010. The absence of any excess above Standard Model (SM) predictions allows one to already set interesting constraints on BSM physics, and several groups of authors have already studied the impact of these results on Supersymmetric scenarios (see, e.g., \\cite{Buchmueller:2011aa,Allanach:2011ut,Akula:2011dd,arXiv:1104.3572, Strumia:2011dv,Buchmueller:2011ki,Bertone:2011nj}).  The prospects for discovering BSM physics at the LHC and the consequences for Dark Matter searches have been thoroughly discussed in the literature. In particular, we have recently discussed the case of positive detection both at the LHC {\\it and} with Dark Matter experiments \\cite{Bertone:2010rv}. Here, we focus instead on the unfortunate case where the LHC {\\it fails} to discover BSM physics, a scenario sometimes referred to as the ``Nightmare scenario'' of particle physics (see, e.g., \\cite{nature} and references therein). We focus on the constrained Minimal Supersymmetric Standard Model (cMSSM) \\cite{CTP-TAMU-24-92,RAL-92-005,hep-ph/9311269,kkrw94,sugra-reviews}, for which detailed studies exist regarding the reach of the LHC for various configurations of beam energy and total integrated luminosity.  The main focus of this paper is to analyse the consequences of the nightmare scenario for Dark Matter searches. These can be actually divided into two broad classes, viz. {\\it direct} and {\\it indirect} searches (for recent reviews see, e.g., \\cite{Bertone:2010zz, Bergstrom,Munoz, Bertone:2004pz,Feng:2010gw}). Direct searches are based on the search for rare nuclear recoils due to the scattering of Dark Matter particles off nuclei in underground experiments. There are many undergoing and upcoming experiments (see, e.g., \\cite{Pato:2010zk} for a discussion of upcoming experimental capabilities), and despite some intriguing signals that have been discovered by the DAMA/LIBRA \\cite{dama} and, more recently, the CoGeNT \\cite{cogent,Aalseth:2011wp,Hooper:2010uy} collaborations, it appears difficult to reconcile a standard Dark Matter interpretation of these results with other experimental findings, in particular with the null searches from XENON-100 \\cite{xe100} or CDMS II \\cite{cdms10} (see, e.g., the discussions in \\cite{gelmini1,Bezrukov,Arina:2011si,Schwetz:2011xm}).  Indirect searches are based instead on the astrophysical searches for secondary particles produced in the annihilation or decay of Dark Matter particles. Despite the huge interest attracted by astrophysical data (see, in particular, the plethora of literature devoted to explaining the positron excess recently measured by PAMELA  \\cite{Adriani:2008zr}), a convincing identification of Dark Matter based on these data seems problematic, due to the lack of strong constraints on the properties of Dark Matter particles as well as to the poor control of astrophysical backgrounds and associated systematics (see, e.g., the discussion in \\cite{nature}). Among the possible smoking-gun signals that can be obtained with indirect searches, we will focus on one of the most intriguing, and less affected by astrophysical uncertainties: the detection of high-energy neutrinos from the centre of the Sun \\cite{Silk:1985ax} using the IceCube neutrino telescope, which is located at the geographical South Pole (see, e.g., \\cite{Halzen:2009vu}).  We show here that upcoming ton-scale direct detection (DD) experiments, and to a lesser extent IceCube, can actually probe a large portion of cMSSM parameter space that will remain unconstrained in the nightmare scenario, and therefore represent {\\it a unique opportunity to discover new physics by the end of the decade}, in case of null searches at the LHC.  The paper is organized as follows: in Sec.\\,\\ref{sec:cmssm} we introduce our theoretical framework, i.e., the cMSSM,  and we present the statistical tools adopted to identify the region of cMSSM parameter space corresponding to the LHC {\\it nightmare} scenario; in Sec.\\,\\ref{sec:results} we present our results and discuss the consequences for Dark Matter searches; and, finally, in Sec.\\,\\ref{sec:summary} we provide a brief summary.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:summary} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  In this study we have quantitatively assessed the potential for discovering the cMSSM in case of null searches at the LHC (the ``nightmare scenario\" of particle physics), with future direct detection experiments and with the IceCube neutrino telescope.  Our main conclusion is that Phase 3 direct detection experiments (that can reach cross sections down to $\\sigsip\\sim10^{-11}$\\,pb) will be able to probe entirely the favoured region of the cMSSM parameter space in the nightmare scenario of particle physics, therefore providing a unique opportunity to test SUSY even in case of null searches at the LHC. Interestingly, these experiments are expected to be built on a timescale of 5--10 years, similar to that required for the LHC to reach IL\\,=\\,100\\,fb$^{-1}$ of data at $\\sqrt{s}=$14\\,TeV.  In our analysis, we fixed all astrophysical parameters to the standard values commonly adopted in the literature. In principle, however, one should take into account all particle physics and astrophysical uncertainties in order to combine in a self-consistent fashion accelerator and particle astrophysics experiments \\cite{Bertone:2010rv,Pato:2010zk, Serpico:2010ae}. Including the uncertainty on the local DM density \\cite{Pato:2010yq, Salucci:2010qr} would not change the fraction of cMSSM points accessible to direct detection experiments, since the SUSY points and the sensitivity curves would be rescaled by the same quantity (i.e., the ratio of the `true' local density over the value adopted here). A more careful treatment of the velocity distribution and of the hadronic uncertainties in the neutralino-nucleon cross-section may instead have a stronger impact on direct and indirect searches \\cite{Ellis:2008hf, Strigari:2009zb,Serpico:2010ae}, but given the ample margin between the bulk of the nightmare cMSSM parameter space and the sensitivity of Phase 3 direct detection experiments, it is unlikely that our main conclusion would change significantly.  Finally we have studied the implications of Higgs searches at the LHC, and we have shown that if a Standard Model-like Higgs is not found at the LHC with IL\\,=\\,100\\,fb$^{-1}$ of data at $\\sqrt{s}=$14\\,TeV, that would basically rule out the cMSSM, while an actual detection in the appropriate mass range would provide additional motivation to continue the study of Supersymmetry with astroparticle experiments.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1107/1107.1540_arXiv.txt": {
        "abstract": "{ We study in detail the power spectra of scalar and tensor perturbations generated during inflation in loop quantum cosmology (LQC). After clarifying in a novel quantitative way how inverse-volume corrections arise in inhomogeneous settings, we show that they can generate large running spectral indices, which generally lead to an enhancement of power at large scales. We provide explicit formul\\ae\\ for the scalar/tensor power spectra under the slow-roll approximation, by taking into account corrections of order higher than the runnings. Via a standard analysis, we place observational bounds on the inverse-volume quantum correction $\\delta \\propto a^{-\\sigma}$ ($\\sigma>0$, $a$ is the scale factor) and the slow-roll parameter $\\epsilon_\\V$ for power-law potentials as well as exponential potentials by using the data of WMAP 7yr combined with other observations. We derive the constraints on $\\delta$ for two pivot wavenumbers $k_0$ for several values of $\\delta$. The quadratic potential can be compatible with the data even in the presence of the LQC corrections, but the quartic potential is in tension with observations. We also find that the upper bounds on $\\delta (k_0)$ for given $\\sigma$ and $k_0$ are insensitive to the choice of the inflaton potentials. } ",
        "introduction": "Both the construction of quantum gravity and the question of its observational tests are beset by a host of problems. On the one hand, quantum gravity, in whatever approach, must face many mathematical obstacles before it can be completed to a consistent theory. On the other hand, assuming that a consistent theory of quantum gravity does exist, dimensional arguments suggest that its observational implications are of small importance. In the realm of cosmology, for instance, they are estimated to be of the tiny size of the Planck length divided by the Hubble distance.  Between the two extremes of conceptual inconsistency and observational irrelevance lies a window of opportunity in which quantum gravity is likely to fall. It is true that we do not yet know how to make quantum gravity fully consistent, and it is true that its effects for early-universe cosmology should be expected to be small. But in trying to make some quantum-gravity modified cosmological equations consistent, it has been found that there can be stronger effects than dimensional arguments suggest. Consistency requirements, especially of loop quantum gravity, lead to modified spacetime structures that depart from the usual continuum, implying unexpected effects.  To partially bridge the gap between fundamental developments and loop quantum gravity phenomenology and observations, one considers effective dynamics coming from constraint functions evaluated on a particular background and on a large class of semiclassical states. In generally-covariant systems, the dynamics is fully constrained, and the constraint functionals on phase space generate gauge transformations obeying an algebra that reveals the structure of spacetime deformations. The algebra of these gauge generators thus shows what underlying notion of spacetime covariance is realized, or whether covariance might be broken by quantum effects, making the theory inconsistent. If a consistent version with an unbroken (but perhaps deformed) gauge algebra exists, it can be evaluated for potential observational implications.  These effective constraints and their algebra, at the present stage of developments, can be evaluated for perturbative inhomogeneities around a cosmological background. While loop quantum gravity is background independent in the sense that no spacetime metric is assumed before the theory is quantized, a background and the associated perturbation theory can be introduced at the effective level. In the case of interest here, the background is cosmological, a flat Friedmann--Robertson--Walker (FRW) spacetime with a rolling scalar field and perturbed by linear inhomogeneous fluctuations of the metric-matter degrees of freedom. The idea is to implement as many quantum corrections as possible, and study how the inflationary dynamics is modified. In this context, it is important that no gauge fixing be used before quantization, as such a step would invariably eliminate important consistency conditions by fiat, not by solving them. The result would be a framework whose ``predictions'' depend on how the gauge was fixed in the first place.  This program, challenging as it is, has been carried out only partially so far, and in gradual steps. First, the full set of constraints was derived for vector \\cite{BH1}, tensor \\cite{BH2}, and scalar modes \\cite{BHKS,BHKS2} in the presence of small inverse-volume corrections. The gravitational wave spectra have been studied in \\cite{CMNS2,CH}, where the effect of inverse-volume corrections and their observability were discussed. The scalar inflationary spectra and the full set of linear-order cosmological observables was then derived in \\cite{InflObs}, thus making it possible to place observational bounds on the quantum corrections themselves \\cite{BCT} (see also refs.~\\cite{inflationworks} for early related works). The reason why these studies concentrate on inverse-volume corrections is mainly technical; in fact, the closure of the constraint algebra has been verified only in this case (and only in the limit of small corrections).  A class of consistent constraints with a closed algebra is known also for vector and tensor modes in the presence of holonomy corrections \\cite{BH1,BH2,MCBG}, and since recently also for the scalar sector, where anomaly cancellation is more difficult to work out \\cite{Bar}. Inspections of cosmological holonomy effects have so far been limited to the tensor sector \\cite{CMNS2,Mie1,GB1,GB2,GB3}.  Exactly isotropic minisuperspace models, where the situation is reversed and holonomy corrections are easier to implement than inverse-volume corrections, provide another reason why it is interesting to focus attention on inverse-volume corrections. Effective equations available for certain matter contents with a dominating kinetic energy \\cite{BouncePert,HighDens} suggest that holonomy corrections are significant only in regimes of near-Planckian densities \\cite{APSII} but not during the timespan relevant for early-universe cosmology, including inflation. Inverse-volume corrections, on the other hand, do not directly react to the density but rather to the discreteness scale of quantum gravity, which is not determined immediately by the usual cosmological parameters. The question of whether they are small or can play a significant role must be answered by a self-consistent treatment.  Such a treatment shows that inverse-volume corrections present an example of quantum-gravity effects that can be larger than what dimensional arguments suggest \\cite{BCT}.  Here we present the full details of the analysis briefly reported in \\cite{BCT} for a quadratic inflationary potential, enriching it with new constraints on other potentials. From a cosmological perspective, we shall provide the complete set of slow-roll equations as functions of the potential, extend the likelihood analysis to quartic and exponential potentials, and discuss how the experimental pivot scale and cosmic variance affect the results.  Before examining the details and experimental bounds of the model, from a quantum-gravity perspective we will clarify some conceptual issues which must be taken into account for a consistent treatment of inverse-volume corrections. In particular, we justify for the first time why inverse-volume corrections depend only on triad variables, and not also on connections. Until now, this was regarded as a technical assumption devoid of physical motivations. Here we show it as a consequence of general but precise semiclassical arguments. Further, we spell the reason why inverse-volume corrections are not suppressed at the inflationary density scale.  The plan of the paper is as follows. We begin in section \\ref{invo} by discussing inverse-volume corrections in LQC and their justification in inhomogeneous models, providing the first implementation for a class of semiclassical states sufficiently large to be used in effective equations. The lattice refinement picture has been introduced and developed in a number of papers, but a major twofold open issue remains. On one hand, there is the need to justify why inverse-volume corrections depend only on triad variables and not on holonomies. On the other hand, it is not clear how the fundamental discreteness scale, previously introduced \\emph{in media res}, arises in the lattice refinement framework. Section \\ref{invo} does both systematically for the first time. The relation between (the size of) inverse-volume and holonomy quantum corrections is clarified in section \\ref{23}, while section \\ref{24} is a recapitulation of past criticism and a discussion of how it is addressed by the present arguments. After that, we turn to an application for observational cosmology. In section \\ref{hsr} we review the formul\\ae\\ of the cosmological observables in the Hubble slow-roll tower \\cite{InflObs}. In section \\ref{vsr} we reexpress these quantities in the slow-roll parameters as functions of the inflationary potential. In section \\ref{ksr} the observables are recast as functions of the momentum pivot scale for a ready use in numerical programs. The effect of cosmic variance on the scalar power spectrum is also discussed therein. In section \\ref{likelihoodsec} we shall carry out the likelihood analysis to constrain the inverse-volume corrections in the presence of several different inflaton potentials by using the observational data of Cosmic Microwave Background (CMB) combined with other datasets. ",
        "conclusions": "In the presence of the inverse-volume corrections in LQC, we have provided the explicit forms of the scalar and tensor spectra convenient to confront inflationary models with observations. Even if the LQC corrections are small at the background level, they can significantly affect the runnings of spectral indices. We have consistently included the terms of order higher than the scalar/tensor runnings. Inverse-volume corrections generally lead to an enhancement of the power spectra at large scales.  Using the recent observational data of WMAP 7yr combined with LSS, HST, SN Ia, and BBN, and analyzing them with techniques routinely used also in standard inflation, we have placed constraints on the power-law potentials $V(\\vp)=V_0\\vp^n$ ($n=2, 4$) as well as the exponential potentials $V(\\vp)=V_0 \\rme^{-\\kappa \\lambda \\vp}$. The inflationary observables (the scalar and tensor power spectra ${\\cal P}_{\\rm s}$, ${\\cal P}_{\\rm t}$ and the tensor-to-scalar ratio $r$) can be written in terms of the slow-roll parameter $\\epsilon_\\V=(V_{,\\varphi}/V)^2/(2\\kappa^2)$ and the normalized LQC correction term $\\delta$. We have carried out a likelihood analysis by varying these two parameters as well as other cosmological parameters for two pivot wavenumbers $k_0$ ($0.002$ Mpc$^{-1}$ and $0.05$ Mpc$^{-1}$).  The observational upper bounds on $\\delta (k_0)$ tend to be smaller for larger values of $k_0$. In table  \\ref{tab1} we listed the observational upper limits on $\\delta (k_0)$ as well as the theoretical priors $\\delta_{\\rm max}$ for the quadratic potential $V(\\vp)=V_0\\vp^2$ with a number of different values of the quantum gravity parameter $\\sigma$ (which is related to $\\delta$ as $\\delta \\propto a^{-\\sigma}$). For larger $\\sigma$, we find that $\\delta (k_0)$ needs to be suppressed more strongly to avoid the significant enhancement of the power spectra at large scales. When $\\sigma \\lesssim 0.5$ the observational upper limits of $\\delta (k_0)$ exceed the theoretical prior $\\delta_{\\rm max}$, which means that the expansion in terms of the inverse-volume corrections can be trustable for $\\sigma \\gtrsim 0.5$.  As we see in Figs.~\\ref{en2}-\\ref{expfig} and in table  \\ref{tab1}, the observational upper bounds on $\\delta (k_0)$ for given $k_0$ and $\\sigma$ are practically independent of the choice of the inflaton potentials. This property comes from the fact that the LQC correction for the wavenumber $k$ is approximately given by $\\delta (k)=\\delta (k_0) (k_0/k)^{\\sigma}$, which only depends on $k_0$ and $\\sigma$. On the other hand the constraints on the slow-roll parameter $\\epsilon_{\\V}$ are different depending on the choice of the inflaton potentials. We have found that the quadratic potential is consistent with the current observational data even in the presence of the LQC corrections, but the quartic potential is under an observational pressure. For the exponential potentials the larger values of $\\epsilon_{\\V}$ are favored compared to the power-law potentials. However, the exponential potentials are not regarded as a realistic scenario unless there is a graceful exit from inflation.  The exponential term $e^{-\\sigma x}=(k_0/k)^\\sigma$ in eq.~\\Eq{Psfinal} is responsible for the enhancement of the power spectrum at large scales. This feature is characteristic of the model and is not reproduced by other sources. For instance, non-commutative geometry or string corrections \\cite{Maartens} predict a suppression, rather than an enhancement, of the spectra. Also some old papers on LQC advertized a suppression of power (e.g., the second reference of \\cite{inflationworks}), but the quantum corrections were not under full control at the level of perturbation theory; the present results supersede those early discussions. Moreover, the signatures of the LQC spectra cannot be mimicked by any standard scalar potential. The enhancement of power is due to a scale-dependent correction in the spectral amplitudes, while exotic potentials would not affect the perturbed dynamical equations (compare with the LQC Mukhanov equations \\cite{InflObs}). In particular, the consistency relation \\Eq{tts2} is notably different with respect to the standard relation $r=-8 n_{\\rm t}$. Finally, also Wheeler--DeWitt quantum cosmology predicts an enhancement of power at large scales \\cite{KiK}, but the effect is qualitatively different from the structure of $\\dpl$ and its size is much smaller. The main reason is that it is governed by the energy scale of inflation, contrary to what happens in LQC.  If we compare the observational upper bounds with the theoretical lower bounds discussed in section \\ref{invo}, we can see that estimates of these parameters are separated by at most a few orders of magnitude, much less than is usually expected for quantum gravity. By accounting for fundamental spacetime effects that go beyond the usual higher-curvature corrections, quantum gravity thus comes much closer to falsifiability than often granted.  It is of interest to see how the future high-precision observations such as PLANCK will constrain the LQC correction as well as the slow-roll parameters. Even in the case where the quadratic potential were not favored in future observations, it would be possible that the small-field inflationary models be consistent with the data.  For these general inflaton potentials, the effect of inverse-volume corrections on the CMB anisotropies should be similar to that studied in this paper."
    },
    "1107/1107.4962_arXiv.txt": {
        "abstract": "Recent cosmological observations, including measurements of the CMB anisotropy and the primordial helium abundance, indicate the existence of an extra radiation component in the Universe beyond the standard three neutrino species. In this paper we explore the possibility that the extra radiation has isocurvatrue fluctuations. A general formalism to evaluate isocurvature perturbations in the extra radiation is provided in the mixed inflaton-curvaton system, where the extra radiation is produced by the decay of both scalar fields. We also derive constraints on the abundance of the extra radiation and the amount of its isocurvature perturbation. Current observational data favors the existence of an extra radiation component, but does not indicate its having isocurvature perturbation. These constraints are applied to some particle physics motivated models. If future observations detect isocurvature perturbations in the extra radiation, it will give us a hint to the origin of the extra radiation. ",
        "introduction": "It is known that most of the energy content of the present Universe consist of ``dark'' components: dark matter and dark energy. It is believed that the remaining components are well-known stuff: baryons, photons and neutrinos. These are ingredients of the standard cosmology supported by cosmological observations.  However, there is no reason to exclude the existence of additional component other than listed above as long as its abundance is not large so that the cosmological evolution scenario is not much affected. Actually, the WMAP measurement of the cosmic microwave background (CMB) anisotropy combined with observations of the baryon acoustic oscillation (BAO) and the Hubble parameter (H0) constrains the effective number of neutrino species $N_{\\rm eff} $ as $N_{\\rm eff}=4.34_{-0.88}^{+0.86}$ (68\\%C.L.) \\cite{Komatsu:2010fb}. Including the Atacama Cosmology Telescope (ACT) data, the constraint is slightly improved as $N_{\\rm eff}=4.56 \\pm 0.75$ (68\\%C.L.)~\\cite{Dunkley:2010ge}. The recent results from South Pole Telescope (SPT), combined with WMAP, BAO and H0 indicate $N_{\\rm eff} = 3.86\\pm 0.42$ (68\\%C.L.)~\\cite{Keisler:2011aw}. Thus these observations suggest the presence of an extra radiation component beyond the standard three neutrino species at the $2\\sigma$ level \\footnote{ There are several studies which discuss how significant the deviation ($N_\\mathrm{eff}\\ne3.04$) is. For detailed discussion, we refer to Refs.~\\cite{GonzalezMorales:2011ty,Hamann:2011hu} }. (See also Refs.~\\cite{Mangano:2006ur,Cirelli:2006kt,Ichikawa:2008pz,Simha:2008zj} for former analyses.) Hereafter we call non-interacting relativistic energy component other than neutrinos as ``extra radiation''.  The CMB measurement is sensitive to the expansion rate of the Universe, or the total energy density of the Universe, at around the recombination epoch. On the other hand, the big-bang nucleosynthesis (BBN) also gives a constraint on the expansion rate at the cosmic temperature of $T\\sim 1$MeV, where the weak interaction is frozen and the neutron-proton ratio is fixed. Thus the observation of the primordial helium abundance gives a constraint on the extra radiation component at $T\\sim 1$MeV. Recently, it was reported that the primordial helium abundance derived from the observations of the extragalactic HII regions indicates an excess at the $2\\sigma$ level compared with the theoretical expectation using the baryon number obtained from the WMAP data~\\cite{Izotov:2010ca}. (See also Ref.~\\cite{Aver:2010wq}.) Ref.~\\cite{Hamann:2010bk} studied constraints on the effective number of neutrinos $N_{\\rm eff}$ and the mass of extra radiation component using CMB (including WMAP 7yr, ACBAR, BICEP and QUaD), SDSS and H0 data, and similarly claimed the existence of an extra radiation and that its mass should be smaller than $\\sim 0.4$~eV. If these observations are true,  we need a theory beyond the standard model in which a new light particle species exists.\\footnote{% Discrepancy between the prediction and observation of the primordial helium abundance may be solved in the large lepton asymmetry scenario~\\cite{Hansen:2001hi,Kawasaki:2002hq,Popa:2008tb,Mangano:2010ei}. We do not pursue this issue in this paper. } Some cosmological scenarios and particle physics models are proposed in order to explain $\\Delta N_{\\rm eff} \\simeq 1$~\\cite{Ichikawa:2007jv,Krauss:2010xg,Nakayama:2010vs,Fischler:2010xz,deHolanda:2010am}.   The most literature treating the extra radiation so far (implicitly) assumes that the extra radiation has the adiabatic perturbation. However, the property of the fluctuation of the extra radiation depends on its production mechanism. For example, there is a scenario that the decay of a scalar field nonthermally produces an extra radiation component~\\cite{Ichikawa:2007jv,Fischler:2010xz}. In this case the extra radiation has the fluctuation of the source scalar field, and if the scalar field is light during inflation, it obtains large-scale quantum fluctuations. Therefore, the extra radiation can have an independent perturbation from the adiabatic one: % it is an isocurvature perturbation. % Investigating the isocurvature perturbation in the extra radiation may have a potential to distinguish the model of the extra radiation, if future observations further confirm $\\Delta N_{\\rm eff} \\simeq 1$.  Therefore, in this paper we study the effect of the isocurvature perturbation in the extra radiation and derive constraints on them using the currently available datasets. The effect on the CMB anisotropy is similar to that induced by the neutrino isocurvaure perturbation. But we are mainly interested in the case where the energy density of the extra radiation itself causes significant change in the Hubble expansion rate. These effects of the extra radiation with isocurvature perturbation are phenomenologically taken into account by scanning both the effective number of neutrino species $N_{\\rm eff}$ and the neutrino isocurvature perturbation $S_\\nu$.   First we develop formalism for calculating the primordial isocurvature perturbation in the extra radiation in Sec.~\\ref{sec:formalism}. For concreteness we restrict ourselves to the two-scalar field case, which generalizes the curvaton model~\\cite{Lyth:2001nq}. The extra radiation is assumed to be produced from decay of either or both of the scalar fields. We also point out that a large non-Gaussianity may exist in the isocurvature perturbation of the extra radiation. A similar study for the case of CDM/baryon isocurvature perturbations, including their non-Gaussianities, was done in Refs.~\\cite{Kawasaki:2008sn,Kawasaki:2008jy,Langlois:2008vk, Kawasaki:2008pa,Hikage:2008sk,Kawakami:2009iu,Hikage:2009rt, Nakayama:2009cr,Langlois:2010dz,Langlois:2010fe,Langlois:2011hn}. In Sec.~\\ref{sec:signature}, we discuss the observational signatures of our model. In particular, we focus on the CMB angular power spectrum, which would be the best probe of our model for the time being. Then we derive observational constraints on such models with extra radiations and isocurvature perturbations using currently available datasets including WMAP, ACT, BAO and $H_0$ measurement in Sec.~\\ref{sec:constraint}. They are applied to some particle physics-motivated models in Sec.~\\ref{sec:model}. We conclude in Sec.~\\ref{sec:conc}.  We note that author in Ref.~\\cite{Hu:1998gy} also considered a specific model for isocurvature perturbations in extra radiation in the different context. ",
        "conclusions": "\\label{sec:conc}  In this paper, we have investigated isocurvature perturbations in an extra radiation component. We have firstly formulated how primordial isocurvature perturbations between the extra radiation and the Standard Model radiation can be generated from fluctuations of two scalar fields in the inflationary Universe. We gave a detailed description of what isocurvature perturbations in the dark radiation are generated. Our derivation of the isocuravature perturbations is based on the $\\delta N$ formalism and fully non-linear. We have also pointed out that non-Gaussianities can  be generated in such isocurvature perturbations, which are to be investigated in more detail in future works. It would be straightforward to extend our formulation for cases with three or more fields. We have also discussed observational signatures of the isocurvature perturbations in the extra radiation. We have pointed out that our model leads to distinct features in the CMB power spectrum which are different from those predicted by the ordinary curvature and CDM/baryon isocurvature perturbations. We have shown that that our model can be constrained from current cosmological observations. Roughly speaking, CMB combined with BAO and the direct measurement of $H_0$ gives constraints on the abundance of the extra radiation and the amplitude of its isocurvature perturbations as $3\\le N_\\mathrm{eff}\\le5$ and  $\\alpha\\le0.1~(0.01)$ for the uncorrelated (totally correlated/anti-correlated) case. These are also translated into constraints on some limiting scenarios which can be realized in models based on particle physics such as SUSY axion model.    Our model will be tested more precisely by cosmological observations which are ongoing or projected in the near future. In particular, CMB power spectrum from the Planck survey is expected to improve constraints on both $N_\\mathrm{eff}$ and $\\alpha$ by an order of magnitude. Furthermore, non-Gaussianities and other observational signatures  would be also informative."
    },
    "1107/1107.1892_arXiv.txt": {
        "abstract": "We derive field equations of Gauss-Bonnet gravity in $4$ dimensions after dimensional reduction of the action and demonstrate that in this scenario Vainshtein mechanism operates in the flat spherically symmetric background. We show that inside this Vainshtein sphere the fifth force is negligibly small compared to the gravitational force. We also investigate stability of the spherically symmetric solution, clarify the vocabulary used in the literature about the hyperbolicity of the equation and the ghost-Laplacian stability conditions. We find superluminal behavior of the perturbation of the field in the radial direction. However, because of the presence of the non linear terms, the structure of the space-time is modified and as a result the field does not propagate in the Minkowski metric but rather in an \"aether\" composed by the scalar field $\\pi(r)$. We thereby demonstrate that the superluminal behavior does not create time paradoxes thank to the absence of Causal Closed Curves. We also derive the stability conditions for Friedmann Universe in context with scalar and tensor perturbations and we studied the cosmology of the model. ",
        "introduction": "One of the most mysterious discoveries of modern cosmology is related to late time acceleration of universe \\cite{Riess:1998cb,Perlmutter:1998np}. For past one decade or so, extensive efforts have been made to understand the underlying reason of this phenomenon (for a recent review see \\cite{Li:2011sd}). According to the standard lore , the acceleration is caused by an exotic form of matter with large negative pressure called dark energy. The cosmological constant and a variety of scalar field systems as representative of dark energy are consistent with observations. Though some of the scalar field systems with generic features, such as tracker solutions, are attractive but nevertheless they do not address the conceptual problems associated with cosmological constant.  There exist an alternative thinking which advocates the need for a paradigm shift, namely, that gravity is modified at large scales which might give rise to late time acceleration \\cite{Clifton:2011jh}. We know that gravity is modified at small scales and thus it is quite plausible that modification also cures at large distance where it has not been varied directly. As for the small scale corrections, no deviations from Einstein's theory have been yet detected; perhaps we need to probe still higher energies to observe these effects. However, situation is quite different challenging at large scales. Indeed, Einstein's theory is consistent with observations to a high accuracy at solar scales thereby telling us that any modification to gravity should be confronted with tough constraints posed by solar physics. Secondly, as for the late time acceleration, the modification should also be distinguished from cosmological constant.  The aforementioned requirements are difficult to satisfy consistently in f(R) theories of gravity. These theories are equivalent to GR plus a scalar degree of freedom whose potential is uniquely constructed from space time  curvature $R$ \\cite{DeFelice:2010aj}. The scalar degree of freedom should mimic quintessence and should hide its detection at small scales {\\it \\`a la chameleon} \\cite{Khoury:2003aq}.   The investigations show that the generic $f(R)$ models either reduce to GR plus cosmological constant \\cite{Thongkool:2009js} or has a $\\phi$MDE instead of a standard matter epoch \\cite{Amendola:2006we} or hit curvature singularity \\cite{Frolov:2008uf} or give rise an ugly fine tuning \\cite{Faraoni:2011pm} which is the price one has to pay for implementing the chameleon mechanism. One might try to remedy these theories by complementing them by quadratic curvature corrections \\cite{Appleby:2009uf}. However, as demonstrated by Lam Hui et al \\cite{Hui:2009kc}, theories based upon chameleon mechanism lead to a violation of equivalence principal of the order of one.  In view of the aforesaid, we need to look for an alternative mechanism of mass screening. The Vainshtein mechanism \\cite{Vainshtein:1972sx} is the one which gives rise to mass screening: Any modification of gravity in the neighborhood of a massive body within a radius dubbed Vainshtein radius are switched off kinematically. The mechanism was invented to address the discontinuity problem \\cite{vanDam:1970vg,Zakharov:1970cc} of massive gravity {\\it \\`a la } Pauli-Fierz \\cite{Fierz:1939ix}. In this theory, the zero helicity graviton mode $\\phi$  is coupled to the stress of energy momentum tensor gives rise to serious violations of local gravity constraints in the limit of vanishing mass of graviton. Vainshtein pointed out that non-linear effects become crucial in this case. The non-linear derivative term added to Pauli-Fierz Lagrangian was shown to implement the mass screening thereby removing the problem of discontinuity.  It is interesting to note that the non-linear term naturally arises in DGP model \\cite{Dvali:2000hr} in the decoupling limit \\cite{Luty:2003vm} such that Vainshtein mechanism is in built in the theory. The scalar degree of freedom obeys Galilean symmetry in flat space time and is free from ghosts is called {\\it Galileon} \\cite{Nicolis:2008in}. There exist higher order Galileon Lagrangians in flat and curved space-time \\cite{Deffayet:2009wt}. The higher order Lagrangian are necessary for realizing the late time de-Siter solution \\cite{Gannouji:2010au}. The Galileon model appears as a particluar case of the Horndeski action defined in 1974 \\cite{Horndeski}. The action has gained some new attention this year, see \\cite{Deffayet:2009mn} for more details.  We should also mention that Galileons are deeply related to massive gravity \\cite{deRham:2010ik} and Dirac-Born-Infeld systems \\cite{deRham:2010eu,RenauxPetel:2011dv} in the framework of higher dimensional theories \\cite{Hinterbichler:2010xn,VanAcoleyen:2011mj,Goon:2011qf,Burrage:2011bt}. The spherically symmetric solution is studied in \\cite{Kaloper:2011qc}.  In this paper, we consider dimensional reduction of Gauss-Bonnet gravity \\cite{Amendola:2005cr} and look for the possibility of mass screening in the model. We also investigate causal structure of flat spherically symmetric and homogeneous and isotropic backgrounds. ",
        "conclusions": "In this paper we have studied simple extension of General Relativity in the context of Lovelock theory. We derived the equations of the reduced theory in $4$ dimensions. We have shown that locally in a flat spherically symmetric background,  the non linear terms coming from the Gauss-Bonnet term in higher dimensions induce Vainshtein mechanism. We found that in $4+2$-dimensions, the Vainshtein radius can be approximated by, $R_V^4\\simeq \\alpha r_s^2$ whereas in case of $4+N$-dimensions by $R_V^3\\simeq \\alpha r_s$ with $N\\neq 2$. We have shown that at distances lower than the Vainshtein radius, the fifth force is negligibly small compared to the gravitational force.\\\\ We have investigated the behavior of the scalar field inside the Vainshtein sphere and  derived stability conditions of the model. We have reaffirmed that the hyperbolicity of the equations is equivalent to the ghost condition and the stability of the Laplacian. We  found that the model has a superluminal propagation of the perturbation of the scalar field in the flat spherically symmetric solution. This faster than light solution appears as soon as the non linear terms of the model become dominant. We have shown, in the special case of $N=2$ that the causality structure of the space time is well defined even in the presence of the superluminal propagation. In fact, we shown that the field propagates in a space-time which is not anymore the Minkowski one but some kind of \"aether\" because of the presence of the background field $\\pi(r)$. This modification of the structure of the space-time is related to the domination of the non linear terms in the Lagrangian. We observed that the light cone gets wider in the aether as compared to the case of Minkowski spacetime provided that the stability conditions hold thereby demonstrating that no CCCs appears even in the presence of the superluminal propagation.\\\\ Finally in the context of an isotropic and homogeneous Universe, we derived Friedmann equations for the field and established the stability conditions in the context of the scalar and tensor perturbations. We show that the matter phase is absent in the  model under consideration.  It will be interesting to investigate the cosmological dynamics and observational constraints on the model under consideration, in the presence of a potential term, in a separate publication. It is also important to investigate the model with general Lovelock terms in simple and non-trivial topology of extra dimensions. We defer this work to our future investigations."
    },
    "1107/1107.1295_arXiv.txt": {
        "abstract": "We study a three stage dark matter and neutrino observatory based on multi-ton two-phase liquid Xe and Ar detectors with sufficiently low backgrounds to be sensitive to WIMP dark matter interaction cross sections down to $10^{-47}~\\n{cm^2}$, and to provide both identification and two independent measurements of the WIMP mass through the use of the two target elements in a 5:1 mass ratio, giving an expected similarity of event numbers. The same detection systems will also allow measurement of the pp solar neutrino spectrum, the neutrino flux and temperature from a Galactic supernova, and neutrinoless double beta decay of $^{136}$Xe to the lifetime level of $10^{27} - 10^{28}$ y corresponding to the Majorana mass predicted from current neutrino oscillation data. The proposed scheme would be operated in three \\textbf{G}eneric stages G2, G3, G4, beginning with fiducial masses 1-ton Xe + 5-ton Ar (G2), progressing to 10-ton Xe + 50-ton Ar (G3) then, dependent on results and performance of the latter, expandable to 100-ton Xe + 500-ton Ar (G4). This method of scale-up offers the advantage of utilizing the Ar vessel and ancillary systems of one stage for the Xe detector of the succeeding stage, requiring only one new detector vessel at each stage. Simulations show the feasibility of reducing or rejecting all external and internal background levels to a level \\textless 1 events per year for each succeeding mass level, by utilizing an increasing outer thickness of target material as self-shielding. The system would, with increasing mass scale, become increasingly sensitive to annual signal modulation, the agreement of Xe and Ar results confirming the Galactic origin of the signal. Dark matter sensitivities for spin-dependent and inelastic interactions are also included, and we conclude with a discussion of possible further gains from the use of Xe/Ar mixtures. ",
        "introduction": "\\label{sec:overview}  Liquid noble gas detectors have proven capability to identify and distinguish both nuclear recoil and electron recoil events~\\cite{Aprile:2010}, and are shown in this paper to have the potential of achieving the ultra-low backgrounds needed to detect weakly interacting dark matter particles (WIMPs) with nucleon interaction cross sections down to $10^{-47}\\,\\n{cm^{2}}$ and with sufficient spectral precision to estimate their mass. With multi-ton target masses, these detectors could at the same time measure the spectra of pp solar and galactic supernova neutrinos, and detect neutrinoless double beta decay in $^{\\mathrm{136}}$Xe at the lifetime level ($10^{27} - 10^{28}\\,\\n{y}$) predicted for Majorana neutrinos from neutrino oscillation data.  It is shown in Sec.\\ref{sec:sensitivity}, that for unambiguous identification and mass measurement of dark matter particles, it is essential to have detectors observing signals in targets with two different atomic numbers, in particular to exploit the $A^2$ dependence of the spin-independent cross section. For this, liquid Xe and liquid Ar provide an ideal pair of targets, having similar operating principles and construction, but having spectral and $A^2$ differences giving approximately a factor 5 difference in event rates. Hence similar signal rates should be achievable for a liquid argon detector having five times the mass of a corresponding liquid xenon detector.  We study a dark matter and neutrino observatory of this type constructed in three progressively larger stages, with \\textbf{G}eneric designs referred to as G2, G3 and G4. The sizes and masses of each stage are shown in Tab.\\ref{tab:g2g3g4_1}. \\begin{table}[!htbp] \\centering \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline & \\multicolumn{2}{c|}{G2} & \\multicolumn{2}{c|}{G3} & \\multicolumn{2}{c|}{G4 \\par } \\\\ \\hline & Xe & Ar & Xe & Ar & Xe & Ar \\\\ \\hline \\multicolumn{1}{|l|}{Target dimensions, masses } & & & & & & \\\\ \\hline \\multicolumn{1}{|r|}{ diameter $\\times$ height (m)} & $1 \\times 1$ & $2 \\times 2$ & $2 \\times 2$ & $4 \\times 4$& $4 \\times 4$ & $8 \\times 8$ \\\\ \\hline \\multicolumn{1}{|r|}{ total target mass (t)} & 2.2 & 9 & 18 & 73 & 146 & 580 \\\\ \\hline \\multicolumn{1}{|r|}{ nominal fiducial target mass (t)} & 1& 5 & 10 & 50 & 100 & 500 \\\\ \\hline \\multicolumn{1}{|l|}{No. of photodetectors } & & & & & & \\\\ \\hline \\multicolumn{1}{|r|}{ top} & 120 (3'') & 600 (3'') & 600 (3'') & 670 (6'') & 670 (6'') & 2000 (6'') \\\\ \\hline \\multicolumn{1}{|r|}{ sides (if instrumented)} & 520 (3'') & 670 (6'') & 670 (6'') & 2400 (6'') & 2400 (6'') & 8000 (6'') \\\\ \\hline \\multicolumn{1}{|r|}{ bottom} & 120 (3'') & 160 (6'') & 160 (6'') & 670 (6'') & 670 (6'') & 2000 (6'') \\\\ \\hline \\end{tabular} \\caption{Summary of sizes, masses and photodetector numbers for G2, G3 and G4 detectors} \\label{tab:g2g3g4_1} \\end{table} \\begin{figure}[!htbp] \\centering \\includegraphics[width=0.9 \\columnwidth]{g2inner.pdf} \\caption{Main parameters of the $1~\\n{ton}/5~\\n{ton}$ (fiducial) G2 system} \\label{fig:g2inner} \\end{figure} \\begin{figure}[!htbp] \\centering \\includegraphics[width=0.9 \\columnwidth]{g3inner.pdf} \\caption{Main parameters of the $10~\\n{ton}/50~\\n{ton}$ (fiducial) G3 system} \\label{fig:g3inner} \\end{figure} \\begin{figure}[!htbp] \\centering \\includegraphics[width=0.9 \\columnwidth]{g2water.pdf} \\caption{G2 system (1t Xe/5t Ar) in water and liquid scintillator shields} \\label{fig:g2water} \\end{figure} \\begin{figure}[!htbp] \\centering \\includegraphics[width=0.9 \\columnwidth]{g3water.pdf} \\caption{G3 system (10t Xe/50t Ar) in water and liquid scintillator shields} \\label{fig:g3water} \\end{figure}  Each stage has the required 5:1 ratio of the Ar and Xe fiducial masses, beginning with 1 ton of Xe and 5 tons of Ar. After a period of data-taking with the G2 system, subsequent scale up can then be achieved by the following steps: \\begin{pbenumerate} \\item Replacing the 5-ton Ar with an equal volume of Xe, taking advantage of the factor 2 density ratio to give a G3 Xe target an order of magnitude larger in mass than G2 Xe. \\item Construction of one new vessel - a matching G3 Ar target, which would again have a mass 5 times that of the G3 Xe mass. \\item If warranted by the performance of the G3 system, replacing the G3 argon by 100 tons xenon to give the order of magnitude G4 scale up. \\item A final option of constructing a G4 Ar target with 5 times the G4 Xe fiducial mass. \\end{pbenumerate}  The detectors will operate using the two-phase principle~\\cite{Dolgoshein:1970}, yielding both scintillation and ionization signals whose ratio provides a factor $\\sim 10^{2} - 10^{3}$ discrimination between nuclear recoils and background electron recoils. This has been demonstrated at the target mass level of $10-20$ kg, both technically and operationally, by three collaborations ZEPLIN~\\cite{Alner:2007,Lebedenko:2009,Lebedenko_1:2009,Akimov:2011}, XENON~\\cite{Angle:2008,Sorensen:2010,Aprile_1:2010} and WARP~\\cite{Benetti:2008}, the latter using liquid Ar and the others liquid Xe. Such devices have demonstrated stable operation for periods in excess of a year. A 60 kg fiducial mass (48 kg fiducial mass) detector in the XENON series is now in successful operation~\\cite{Aprile_1:2010}. In this paper we study the achievable backgrounds and potential physics sensitivity of the detection system in Tab.\\ref{tab:g2g3g4_1}, without further consideration of hardware design.  Illustrations of the basic internal structure of the G2 and G3 systems are shown in Fig.\\ref{fig:g2inner} and Fig.\\ref{fig:g3inner}. Fig.\\ref{fig:g2water} and Fig.\\ref{fig:g3water} show the detectors immersed in water and liquid scintillator shields (to be discussed further in \\ref{app:externalbkg}). A constant water-shielding diameter is used to illustrate both systems, although smaller water shield thicknesses would be acceptable for the G2 system. The G4 Xe detector would be identical in configuration to the G3 Ar detector of Fig.\\ref{fig:g3inner}. The photodetectors shown in Fig.\\ref{fig:g2inner} and Fig.\\ref{fig:g3inner} register both the direct scintillation light from each event (referred to as S1) and also the simultaneous ionization which is drifted to the liquid surface by an electric field and then extracted and further drifted to produce proportional scintillation in the gas (referred to as S2). The electric drift field will be typically 0.5 - 1 kV/cm. XENON100~\\cite{Aprile:2012} operated at 0.5 kV/cm, while XENON1T~\\cite{Aprile:dm2012} is designed at 1 kV/cm with a drift length of 100 cm, for which the total voltage 100 kV has been achieved in the R\\&D phase of that project. Thus the voltage requirements of the G2 stage are already demonstrated. The G3 stage would need a voltage of 200 kV for the xenon detector, which will also suffice for the argon detector at a field of 0.5 kV/cm, and a liquid argon TPC with this maximum voltage is already under development for the LBNE experiment~\\cite{LBNE}. A G4 xenon detector would require either a (foreseeable) factor 2 improvement over the LBNE value or a compensating reduction in the drift field.  The ratio S2/S1 provides a parameter which separates nuclear recoil and electron recoil events into two populations with less than 1\\% overlap. For the photodetectors we envisage the use of 3 inch or 6 inch QUPIDs (avalanche photodiodes within a quartz enclosure~\\cite{Arisaka:2009}), which are designed to have a factor 100 lower U/Th activity, per unit area of coverage, than the best photomultipliers. The required numbers of these are summarised in the lower half of Tab.\\ref{tab:g2g3g4_1}. The QUPIDs have a quantum efficiency $\\sim 35-37\\%$~\\cite{Teymourian:2011} for the 180 nm Xe wavelength, while for the liquid argon detectors, a TPB coating~\\cite{TPBref} would be used to wavelength-shift the argon scintillation light into the QUPID range, again with a QE $\\sim 35-37\\%$~\\cite{TPBAr:UCLA}. Development studies~\\cite{Teymourian:2011} show that the QUPID dynamic range is sufficient to cover both the low energy dark matter searches and the MeV-range neutrinoless double beta decay search.  Ideally the QUPIDs would fully surround the detector in a $4\\pi$ array, but a lower-cost option (at some sacrifice of light collection, energy threshold, and position resolution) would be to locate the QUPIDs at top and bottom only, using high-reflectivity ($>95\\%$) PTFE on an array of $10-12$ flat panels around the sides of the target volume (illustrated for the G2 Xe detector in Fig.\\ref{fig:g2inner}). Simulations below will show that, despite the greater simplicity and lower cost of these reflecting panels, the use of QUPIDs on the side walls would improve the light collection, and hence the energy threshold and detector sensitivity, by a factor 2 throughout most of the target volume.  In Sec.\\ref{sec:sensitivity} we use the basic expressions for WIMP nuclear recoil spectrum (on the assumption of a canonical velocity distribution), and spin-independence of the WIMP-nucleon cross section, to determine the sensitivity of each of these detector systems to WIMP-nucleus cross-section and WIMP mass, on the assumption that backgrounds (from neutron, gamma, and electron recoil events) can be reduced to the level $\\sim~0.1-0.5$ background events per year -- i.e. signals or limits not background limited. We also show the sensitivity to the expected annual modulation arising from the Earth's orbital and Galactic motion, and the sensitivity to other types of WIMP-nucleus interaction, in particular spin-dependent and hypothetical inelastic interactions.  In Sec.\\ref{sec:backgrounds} we discuss in detail simulation results of both external water shielding and the reduction of detector backgrounds from local radioactivity, in the photodetectors and detector vessel, showing that a combination of a liquid scintillator veto and self-shielding by an outer thickness of target material provides the key to achieving an essentially background-free fiducial target volume in all three of the G2, G3 and G4 systems. In addition we discuss the reduction of backgrounds from radioactive contaminants of liquid Xe and Ar.  In Sec.\\ref{sec:nulowenergy} we calculate the sensitivity of these detectors to neutrinos from the solar pp chain and from a Galactic supernova. Finally in Sec.\\ref{sec:nulessdoublebeta} we show that sufficiently low backgrounds could be achieved in the G3 or G4 systems for the observation of neutrinoless double beta decay from $^{136}$Xe at the lifetime level 10$^{27}$ $-10^{28}~\\n{y}$ corresponding to the Majorana mass estimated from neutrino mixing data. Other possible configurations for this include a concentric target of enriched $^{136}$Xe shielded by $^{136}$Xe-depleted Xe (the ``XAX'' scheme, described in ~\\cite{Arisaka:2009}) and the use of Xe/Ar mixtures. ",
        "conclusions": "\\label{sec:conclusion}  We have studied the capabilities of a several stage program of construction and operation of multi-ton liquid Xe and liquid Ar two-phase TPC detectors. The principal objective is to measure the energy spectrum of nuclear recoils from weakly interacting massive particles that may constitute the Galactic dark matter, with sufficiently low background to reach WIMPnucleon cross-sections $\\sim 10^{-46}~\\n{cm^2}$ and lower, with sufficient events to allow determination of the shape of the recoil energy spectrum and from this to estimate the WIMP mass. As previously proposed~\\cite{Smith:2005} it is important to measure this signal in two detectors with different target elements, to confirm the expected $A^2$-dependence of the signal from a coherent spin-independent cross-section, and to provide agreement between two independent determinations of the incident particle mass.  The first stage of this program, referred to as G2, would consist of a liquid Xe detector of 1-ton fiducial mass, paired with a liquid Ar detector with a 5-ton fiducial mass. The second stage, referred to as G3, would use the G2 argon vessel as a 10-ton fiducial Xe detector, then paired with a 50-ton fiducial Ar detector. A further order of magnitude scale-up, referred to as G4, could be achieved by using the G3 argon vessel for a 100 ton Xe detector, and adding if required, a 500 ton argon detector. It has been shown in this paper that sufficiently low gamma and neutron backgrounds can be achieved by combining the following methods: \\begin{pbitemize} \\item[a)] by the now proven two-phase discrimination of nuclear recoils from gamma/beta events; \\item[b)] by using in each case a total target mass approximately twice the fiducial mass and using the outer half of the mass as self-shielding; \\item[c)] by the development, now in progress, of a new high quantum efficiency QUPID photodetector, with a radioactivity level lower by $1-2$ orders than the current best photomultipliers. This would allow for the first time the use of full $4\\pi$ photodetector arrays with improved light collection, energy resolution, and event position resolution. \\end{pbitemize}  Other internal backgrounds, from radioactive impurities in structural materials, have been shown to be reducible to negligible levels. In addition to achieving the best sensitivity for measurement of signals from spin-independent interactions, the data will at the same time provide the best limits on spin-dependent or inelastic WIMP scattering, again with independent data from Xe and Ar detectors.  The two detectors of the G3 system are further capable of observing the annual WIMP-spectrum modulation arising from the combined Earth-Sun motion in the Galaxy. The Xe and Ar detectors will provide two independent measurements of this, confirming the Galactic origin of the signal.  The sensitivity of liquid noble gas detectors to signals below 100 keV provides the opportunity to measure signals from two sources of astrophysical neutrinos: solar and supernova. The excellent position resolution of the two-phase design and light collection system allows the gamma background to be reduced sufficiently to allow pp solar neutrinos to be clearly seen as an neutrino-electron scattering spectrum up to 250 keV and, with a Xe target depleted in $^{136}$Xe, sufficient events could be obtained to measure the spectrum to a precision of a few percent. The total neutrino burst from a Galactic supernova could be observed as a coherent neutral current nuclear recoil spectrum, in both the Ar and Xe detectors, with sufficient accuracy to measure both the total flux and mean energy (temperature) of the neutrino source. At a distance $\\sim 10$ kpc, the precision would be $\\sim 10\\%$ with the G3 detection system.  We have discussed the further gains to be achieved with the G4 100-ton (fiducial) Xe scale-up, with or without a corresponding 500-ton (fiducial) Ar detector. These include order of magnitude gains in dark matter event numbers and mass determination, more precise annual modulation signal, and increased numbers of neutrino events from a supernova burst, giving sensitivity to more remote Galactic supernovae.  An important application of the G3 and/or G4 systems would be a high level of sensitivity to neutrinoless double beta decay from $^{136}$Xe. This would be achievable with the 10-ton G3 Xe detector, using either natural or enriched $^{136}$Xe, reaching a lifetime sensitivity $10^{27} - 10^{28}$ y which, for Majorana neutrinos, lies within the $0.1-0.01$ eV majorana mass range expected from existing neutrino mixing data, and hence an expectation of a positive $0 \\nu \\beta \\beta $ signal for the first time. We consider also the option of using Xe/Ar mixtures for this and/or dark matter detection, for potentially higher resolution and improved self shielding.  This study concludes that the ultra-low backgrounds needed for these searches are achievable with existing techniques, and that combinations of multi-ton Xe and Ar detectors provide the most sensitive method of identifying dark matter signals, providing at the same time new high sensitivity measurements at the frontier of neutrino physics."
    },
    "1107/1107.4686_arXiv.txt": {
        "abstract": "We present the method of describing an evolution in quantum cosmology in the framework of the reduced phase space quantization of loop cosmology. We apply our method  to the flat Friedman-Robertson-Walker model coupled to a massless scalar field.  We identify the physical quantum Hamiltonian that  is positive-definite and generates globally an unitary evolution of considered quantum system. We examine properties  of expectation values of physical observables in the process of the quantum big bounce transition. The dispersion of  evolved observables are studied for the Gaussian state. Calculated relative fluctuations enable an examination of  the semi-classicality conditions and possible occurrence of the cosmic forgetfulness.  Preliminary estimations based on the cosmological data suggest that there was no cosmic amnesia. Presented results are analytical, and numerical computations are only used for the visualization purposes. Our method may be generalized to sophisticated cosmological models including the Bianchi type universes. ",
        "introduction": "The loop quantum cosmology (LQC) method seems to be an efficient method of quantization of cosmological models of general relativity developed recently. Presently, we have two versions of this method: standard LQC (see, e.g. \\cite{Ashtekar:2003hd,Bojowald:2006da,Ashtekar:2006wn} and references therein) and nonstandard LQC \\cite{Dzierzak:2008dy,Dzierzak:2009ip,Malkiewicz:2009qv,Malkiewicz:2009zd, Mielczarek:2010rq,Mielczarek:2010wu,Dzierzak:2009dj,Malkiewicz:2010py}. For an extended {\\it motivation} for developing the nonstandard LQC we recommend an appendix of our paper \\cite{Malkiewicz:2009qv}.   The {\\it standard} LQC has been developed by several groups around the world in the last decade. This approach follows the Dirac program in which one first identifies the kinematical Hilbert space ignoring the dynamical constraints. Next, the dynamical constraint (choosing suitable gauges leads to a single constraint) of the theory is promoted into an operator acting in this space. Kernel of this operator is used to construct the physical Hilbert space. Similarly, observables are analyzed firstly on the kinematical phase space and  suitable unitary transformation is constructed to map the kinematical results into the physical ones.  The {\\it nonstandard} LQC has been proposed recently. In the first stage of this approach one prepares the classical formalism for quantization: (1) Solutions to the Hamilton equations which satisfy the dynamical constraints are found. In the case of complicated system of equations, one may examine the structure of the constraint surface by the phase portrait methods for dynamical systems \\cite{perko}. (2) Elementary Dirac observables on the constraint surface are determined, which define the physical phase space (PFS). (3) Physical observables are defined, i.e. the observables which after quantization can be used for predicting outcomes to be compared with observational data. Physical observables are introduced as functions of elementary Dirac observables and an evolution parameter.  In the second stage, one quantizes the classical system: (1) Self-adjoint representations of physical observables are constructed by using the representation of the algebra of elementary observables. (2) Eigenvalue problems for physical observables are solved to get their spectra. (3) Evolution of expectation values of physical observables is examined. It consists in finding the new Hamiltonian of the classical theory that generates dynamics on the physical phase space. This new Hamiltonian is no longer a dynamical constraint of the classical theory. One finds a self-adjoint representation of this Hamiltonian. Quantum Hamiltonian is used, via Stone's theorem, to define an unitary operator. This operator is used to examine an evolution of the quantum system.  Preliminary results obtained in \\cite{Malkiewicz:2009qv,Malkiewicz:2010py} indicate that proposed method of describing an evolution of a quantum cosmological system is reasonable. In this paper we give {\\it complete} presentation of the evolution. We demonstrate that it is able to reveal the details concerning the nature of the quantum big bounce transition. In particular, we analyze quantum {\\it fluctuations} of physical observables in the propagation across the quantum bounce from the past to the future time infinities.  In the standard LQC an evolution of cosmological system is described quite differently. The kernel equation for the operator constraint is used to construct, after some formal rearrangements, an equation interpreted as  an evolution equation.  Unfortunately, so obtained equation is usually so complicated that one can only solve it by combined analytical and numerical methods.  This is why the preliminary examination of the evolution of the quantum FRW model has shown that classical big bang turns into quantum big bounce transition \\cite{Ashtekar:2006wn}, but could not say anything specific about the nature of the quantum bounce. In particular, an evolution of the dispersion effects of quantum observables was not done satisfactory. Replacing an exact Hamiltonian constraint of the FRW model by a simpler one have enabled making some analytical analysis. This simplified method for describing an evolution, called sLQC \\cite{Ashtekar:2007em}, has shown that the cosmic {\\it amnesia} for the case of {\\it semiclassical} states, discovered earlier in analyzes of a simple cosmological toy model \\cite{BojowaldNature}, does not occur \\cite{Corichi:2007am}.  Our nonstandard LQC method enables analytical studies, with an exact expression for dynamical constraints, of subtle quantum effects of specific cosmological model of  the universe. In this paper, we consider the flat FRW model with a free massless scalar field. The choice of the model results from the fact that the dispersion effects  have been examined so far  mainly for this model. On the other hand,  the model is quite simple and the FRW symmetry is supported by the current observational cosmology.  In order to have our paper self-contained,  we recall in Sec. II A some aspects of  derivation of the Hamiltonian {\\it constraint} of our nonstandard LQC \\cite{Mielczarek:2010rq}. In Sec. II B we introduce the notion of the {\\it physical} Hamiltonian.  In Sec. III we construct the physical {\\it quantum} Hamiltonian and examine its spectral properties. The {\\it evolution} of quantum FRW model is presented in Sec. IV, where we consider the {\\it dispersion} of physical observables. We also briefly evoke the problem of time. The relative {\\it fluctuations} of observables, for the Gaussian state, are considered as a function of time in Sec. V. We conclude in the last section.  Appendix A includes the derivation of the formulas used in the  section on quantum fluctuations. ",
        "conclusions": "In the Hamiltonian formulation of general relativity, GR, the total Hamiltonian is a linear combination of the constraints (see, e.g. Eq. (\\ref{conham})) so it cannot play the role of the generator of an evolution of a gravitational system. On the other hand, the GR system evolves according to the Einstein equations. Can one overcome this difficulty of the Hamiltonian formulation? There are two ways of dealing with this problem: \\begin{enumerate} \\item One eliminates the time variable in favor of any canonical variable by some formal trick carried out on Hamilton's equations (see, e.g. \\cite{Dzierzak:2009ip} for more details), which leads to the  so called {\\it relative} dynamics commonly used in LQC. In this procedure the constraints are used to define the {\\it physical} phase space. However, the evolution is poorly defined since one gets the dependance of canonical variables in terms of any specific variable so the dependance on time may become deeply hidden. \\item Using the constraints, one expresses the specific canonical variable (of Hamilton's equations) in terms of other variables. Elementary Dirac observables are constants of motion and include the dynamical constraint. The new Hamilton's equations (including the constraint) are used for finding the Hamiltonian that generates dynamics on the physical phase space. This new Hamiltonian is no longer a dynamical constraint of the classical theory, but the generator of dynamics, formally {\\it free}  of dynamical constraints. This approach  restores  the notion of an evolution of a classical system. \\end{enumerate}  One should remember that the above considerations on the evolution parameter apply first of all to the situation with `matter' field, like the scalar field, that can be used to define this parameter. It would be interesting to extend these ideas to the case with no sources. The simplest  nontrivial example is the Kasner model.  When we wish to {\\it quantize} a Hamiltonian system with constraints, we may apply the Dirac or the RPS quantization methods. Both methods are  plagued by numerous {\\it ambiguities}. Our analyzes are devoid of the need of any restriction to some superselection sectors that naturally arise in the standard LQC quantization. The latter need not be a drawback of the method as it may leave some interesting imprints in the physical predictions  \\cite{GJT}. The former one, seems to be less complicated than the Dirac method, and offers an {\\it analytical} insight into physical aspect of considered model. Presented results demonstrate, to some extent, that our quantization scheme leads to a quantum {\\it cosmological} system with general properties of a quantum systems we are dealing with in {\\it terrestrial} laboratories.  It seems that one can apply our method to much more complex cosmological models than the FRW universe. Recently, we have managed to quantize the Bianchi I model \\cite{Dzierzak:2009dj,Malkiewicz:2010py}.  The case of the Biachi II model can be treated by analogy. In summary, we suggest that quantization of cosmological systems in terms of the RPS method is much more efficient and unique than Dirac's method. However, since quantum cosmology `experimental' data are not available yet, the best strategy seems to be applying both methods to {\\it compare} the results. An agreement of the results would prove that the procedure of quantization was correct.  The loop quantization methods applied to simple cosmological models teach us that approximating the curvature of connection by holonomies around small loops enable replacing   classical sinularities by quantum bounces \\cite{Ashtekar:2003hd,Bojowald:2006da,Ashtekar:2006wn, Dzierzak:2008dy,Dzierzak:2009ip,Malkiewicz:2009qv,Malkiewicz:2009zd, Mielczarek:2010rq,Mielczarek:2010wu,Dzierzak:2009dj,Malkiewicz:2010py}. To get some information about the nature  of a bounce, one examines propagation   quantum states across the bounce \\cite{Corichi:2007am,BojowaldNature,Bojowald:2007gc,Corichi:2011rt, Kaminski:2010yz}. Such method is similar, to some extent, to the method used, for instance, in nuclear physics where one scatters a particle against an atomic nucleus to get information on the structure of the latter. In papers \\cite{BojowaldNature,Bojowald:2007gc} one considers solvable toy models (motivated by LQC) to argue that a quantum state before the bounce may become decohered at the bounce and become semiclassical afterwards. Applying the sLQC prescription authors of \\cite{Corichi:2007am,Corichi:2011rt} claim that there is no cosmological amnesia at all: suitable semiclassical state before the bounce keeps being semiclassical after the bounce.  Authors of \\cite{Kaminski:2010yz} give strong support to this result by applying various analytical and numerical  methods, within the standard LQC, to general forms of semiclassical states. They identify the condition under which one has the preservation of the semiclassicality across the bounce. However, they have mainly examined  the volume observable  which is little useful for testing the cosmic amnesia.  In our paper we have considered the transition of the Gaussian wave packet across the bounce of the quantum FRW universe within an {\\it exact} framework.   It results from our studies that the $\\hat{Q}$ observable is the proper quantity to study the cosmic amnesia. It is because relative fluctuations of $\\hat{Q}$ can be observationally constrained. Moreover, the semiclassicality condition imposed in the expanding phase restricts also the quantum fluctuations in the contracting phase.  The preliminary observational constraint $\\;\\frac{\\Delta \\hat{Q}}{\\langle \\hat{Q}\\rangle} < 0.02\\;$ indicates  that the semiclassicality condition,  as defined earlier, may be fulfilled. Owing to this, one can infer that there was a cosmic {\\it amnesia} or there was a {\\it recall} depending on what we mean by a semiclassical state. Our results support the prediction of the standard LQC \\cite{Corichi:2007am,Corichi:2011rt,Kaminski:2010yz}. We suggest to repeat the calculations with the  variety of  states different from the Gaussian type  states to verify  our results.  On the other hand, the LQC results obtained for the FRW type models cannot be probably used successfully to describe the Universe. The very high symmetry of space specific to the FRW model is probably unrealistic near the cosmological singularity. We suggest that the real nature of the bounce may become known only after we quantize the Belinskii-Khalatnikov-Lifshitz (BKL) scenario \\cite{Belinski:2009wj,BKL2,BKL3}, which concerns the generic cosmological singularity. Quantization of simple cosmological models carried out during the last decade may be treated as warming up before meeting this challenge.  \\ack We thank Vladimir Belinski, Jean-Pierre Gazeau, Przemys{\\l}aw Ma{\\l}kiewicz and Wies{\\l}aw Pusz  for helpful discussions. JM has been supported by Polish Ministry of Science and Higher Education grant N N203 386437 and by Foundation of Polish Science scholarship START. Also we would like to thank the anonymous referees for the constructive criticisms.  \\appendix"
    },
    "1107/1107.2405_arXiv.txt": {
        "abstract": "We present a new X-ray analysis of the dwarf Seyfert galaxy NGC~4395, based on two archival \\textit{XMM-Newton} and \\textit{Suzaku} observations. This source is well known for a series of remarkable properties: one of the smallest estimated black hole masses among Active Galactic Nuclei (of the order of $\\sim$10$^5 M_{\\sun}$), intense flux variability on very short time-scales (a few tens of seconds), an unusually flat X-ray continuum ($\\Gamma \\sim 1.4$ over the 2--10~keV energy range). NGC~4395 is also characterized by significant variations of the X-ray spectral shape, and here we show that such behaviour can be explained through the partial occultation by circumnuclear cold absorbers with column densities of $\\sim$10$^{22}$--10$^{23}$~cm$^{-2}$. In this scenario, the primary X-ray emission is best reproduced by means of a power law with a standard $\\Gamma \\sim 1.8$ photon index, consistent with both the spectral slope observed at higher energies and the values typical of local AGN. ",
        "introduction": "The central regions of the dwarf spiral galaxy NGC~4395 show most of the common signatures of nuclear activity. The optical and ultraviolet (UV) spectra reveal prominent high-ionization forbidden lines on top of a nearly featureless continuum, and broad wings corresponding to gas velocities in excess of $\\sim$10$^3$~km~s$^{-1}$ are detected in the permitted lines (Filippenko \\& Sargent 1989). Contrary to the objects of the same kind, the emission-line properties, the optical to X-ray variability pattern and the inferred accretion rate of NGC~4395 are those typical of the Seyfert class, of which this source is usually considered to represent the least luminous member. Different methods have been employed in the last years to derive the mass of its central black hole: the estimate obtained through reverberation mapping is $M_\\rmn{BH} \\simeq 3.6 \\times 10^5 M_{\\sun}$ (Peterson et al. 2005), but the lack in this galaxy of a significant bulge and the stringent upper limit of 30~km~s$^{-1}$ on its velocity dispersion suggest an even lower value, of the order of $\\sim$10$^4$--10$^5 M_{\\sun}$ (Filippenko \\& Ho 2003). Anyhow, the engine of NGC~4395 falls somewhere between the stellar-mass black holes found in Galactic X-ray binaries and the supermassive black holes residing inside active galactic nuclei (AGN). As such, it can provide critical information about the relationship between these two populations and the physics of accretion systems in general. \\\\ In the light of all these pieces of observational evidence, NGC~4395 is a true scaled-down version of an ordinary Seyfert galaxy, the only difference with respect to its high-luminosity counterparts being the much smaller mass of the central black hole. On the other hand, the X-ray observations of this source indicate an unusual spectral hardness at $\\sim$2--10~keV, and suggest a wide range of variations for the intrinsic photon index. The most extreme states ($\\Gamma \\sim 0.6$; Moran et al. 2005) are even difficult to interpret within the standard two-phase model (Haardt \\& Maraschi 1991), posing serious questions on the production mechanism of the X-ray emission itself. The presence of undetected absorption effects has been frequently invoked as a possible explanation. Indeed, the stronger flux variability characterizing systematically the soft X-ray bands can be attributed to a complex, multi-zone warm absorber, whose properties have been discussed in detail in several works (Iwasawa et al. 2000; Shih, Iwasawa \\& Fabian 2003; Dewangan et al. 2008). Here we review the two highest-quality observations of NGC~4395, performing in both cases an accurate time-resolved analysis mainly focused on neutral absorption, in order to test whether changes of its column density and/or covering factor play a role in the apparent X-ray spectral hardness of this source. ",
        "conclusions": "We have discussed a possible interpretation of the X-ray spectral hardness usually observed in the low-luminosity active galaxy NGC~4395. This source harbours in its centre a black hole with estimated mass of $\\sim$10$^5 M_{\\sun}$, and it is one of the few known objects whose study can shed light on the links between the physics of accretion processes in Galactic black hole binaries and AGN. In spite of being a genuine Seyfert galaxy in many respects, NGC~4395 remains a somewhat puzzling source because of the inferred flatness of its primary X-ray continuum. The existence of complex absorption effects has often been proposed as a likely explanation. Here we have provided for the first time an example of these effects based on observational evidence, by reviewing the two highest-quality looks of NGC~4395 taken by \\textit{XMM-Newton} and \\textit{Suzaku}: in both cases, a time-resolved analysis shows that the spectral evolution of the source can be interpreted by means of a two-phase neutral absorber with variable covering factor. As a first approximation, this cold absorber is identified with the system of broad-line clouds, allowing for a double-peaked $N_\\rmn{H}$ distribution. The low column density component can otherwise be attributed to an external narrow-line or torus-like region, with nearly constant $f_\\rmn{cov}$. This is presumably an oversimplification of the physical and geometrical structure of the circumnuclear environment (which is known to comprise also a complex, multi-zone warm absorber), but even the existence of a single partial-covering cold screen cannot be completely ruled out on statistical grounds. \\\\ We stress that the absorption variability scenario presented here is not unique, and different models based on intrinsically flat X-ray continua (with significant changes of either the photon index or the reflection strength) can describe the observed behaviour of this source equally well. This interpretation, however, fits into the analogy between the properties of NGC~4395 and those of \\textit{standard} high-luminosity Seyfert galaxies, many of which are systematically affected by X-ray absorption variability due to the clumpiness of their circumnuclear regions. Moreover, it allows us to retrieve in a natural way a $\\Gamma \\simeq 1.8$ photon index below 10~keV, in perfect agreement with the \\textit{Swift}/BAT spectral slope and the usual values measured among AGN. No alternative explanation would then be required for the intrinsic 2--10~keV flatness of NGC~4395."
    },
    "1107/1107.2858_arXiv.txt": {
        "abstract": "With the launch of the $\\textit{Gaia}$ satellite, detection of many different types of transient sources will be possible, with one of them being optical afterglows of gamma-ray bursts (GRBs). Using the knowledge of the satellite's dynamics and properties of GRB optical afterglows, we performed a simulation in order to estimate an average GRB detection rate with $\\textit{Gaia}$. Here, we present the simulation results for two types of GRB optical afterglows, differing in the observer's line of sight compared with a GRB jet axis: regular (on-axis) and orphan afterglows. Results show that for on-axis GRBs, less than 10 detections in five years of foreseen $\\textit{Gaia}$ operational time are expected. The orphan afterglow simulation results are more promising, giving a more optimistic number of several tens of detections in five years. ",
        "introduction": "\\label{s1} European Space Agency's cornerstone mission $\\textit{Gaia}$ is scheduled for launch in 2013. It is primarily an astrometric mission, designed to measure precise 3D positions for up to a billion stars in our Galaxy and other bright objects in the sky. In addition to astrometric measurements, spectroscopy and multi-band photometry will be performed, providing radial velocities and additional astrophysical information on the stars and other objects. The main scientific goal of $\\textit{Gaia}$ mission is to clarify the origin and formation history of our Galaxy.  The goals of the mission are described in \\citet{perryman}, \\citet{lindegren}, and, more thoroughly, in \\citet{turon}.  $\\textit{Gaia}$ will continuously scan the sky, covering the whole sky and building its own catalogue of celestial objects. A preliminary catalogue will be released every 6 months and a final catalogue, exploiting all the data recovered during the mission, will be released at the end of the operational time (around 2020) \\citep{lammers}. Comparing further observations with this catalogue will enable $\\textit{Gaia}$ to detect various transient sources, with some of them being cataclysmic variables, supernovae, active galactic nuclei, asteroids, etc. Since the catalogue of $\\textit{Gaia}$ astrometric and photometric data will be available only late into the $\\textit{Gaia}$ operation, triggers from transient phenomena will actually be the first $\\textit{Gaia}$ data released to the scientific community. In order to be prepared for this stream of data and to make sure the alert stream is accurate, reliable, and free from contamination, the $\\textit{Gaia}$ Science Alerts Working Group has been formed\\footnote{http://www.ast.cam.ac.uk/ioa/research/gsawg}.  One type of possible transients to be detected by $\\textit{Gaia}$ is gamma-ray burst optical afterglows. Gamma-ray bursts (GRBs) are extremely energetic explosions occurring at cosmological distances \\citep{piran,meszaros}. They represent the most luminous events in the gamma-ray part of the electromagnetic spectrum known in the universe. Usually, the short-term prompt gamma-ray emission is followed by an afterglow emitted at longer wavelengths from X-rays to radio waves, which can last for several days. According to the standard fireball model \\citep{reesa, reesb} the prompt gamma emission is produced in the internal shocks of the relativistically expanding ejecta, while the afterglow is synchrotron emission from relativistically accelerated electrons in the external shocks, created in the ejecta interaction with the circumburst medium. In general, afterglow light curves exhibit the power-law decay with time \\citep{sari}, although many additional features like density bumps \\citep{lazzati,guidorzi}, late energy injections \\citep{nousek,evans} and reverse shocks \\citep{akerlofc,fox,kobayashi,gombocb} have been seen in recent years, when better time coverage observations became possible due to accurate GRB alerts from the $\\textit{Swift}$ satellite \\citep{gehrels}  and a number of follow-up robotic telescopes (e.g., the Liverpool Telescope \\citep{gomboc}, ROTSE \\citep{akerlofa}, REM \\citep{zerbi}, etc.).  GRB explosions are believed to be collimated, rather than spherical \\citep{rhoadsa,granot}. An indirect argument supporting this is based on the high values of energy output in gamma rays ($E_{\\rm iso}$), which is of the order of $\\sim M_{\\odot}c^{2}$ if we assume that the explosion is spherically symmetric. It is known that GRBs occur in galaxies, and are of stellar origin (most probably the collapse of massive, rapidly rotating stars and mergers of compact objects). It is difficult to conceive  a model with a stellar progenitor, which can produce such high energies inferred in the case of spherically symmetric explosions. More direct evidence comes from precise observations of afterglow light curves: an achromatic  break in the light curve, which has been observed in many GRB afterglows at a few hours to days after the GRB itself (\\citealt{granot} and references therein), was predicted theoretically prior to observational discovery \\citep{rhoads,sarib}. The break (usually referred to as a jet break) can be well explained by considering a jet with a half-opening angle $\\theta_{\\rm j}$, moving initially at a relativistic speed with high Lorentz factor $\\Gamma \\sim 100 - 1000$. Because of the relativistic beaming effect, the radiation emitted from the jet is beamed in a half-opening angle $\\theta_{\\rm b} = \\frac{1}{\\Gamma}$ in the direction of motion. The expanding jet is surrounded with a medium that causes the jet to decelerate. This results in an increasing value of the angle $\\theta_{\\rm b}$. Shortly after the prompt GRB emission, we cannot distinguish between a spherical and collimated explosion due to strong beaming, since $\\theta_{\\rm j} > \\theta_{\\rm b}$. But after some time, the jet slows down and $\\theta_{\\rm b}$ becomes larger than the opening angle of the jet itself. At this point, the radiation flux drops and we can see an achromatic break in the light curve \\citep{kulkarni,klose}. As mentioned previously, a collimated explosion results in a reduced estimation of energy output, i.e., $E_{\\rm j} = \\frac{1}{2}\\theta_{\\rm j}^{2}E_{\\rm iso}$, where $E_{\\rm j}$ is collimation-corrected energy in the case of a jet with half-opening angle $\\theta_{\\rm j}$.  The shape of the afterglow light curve depends on the relative position of the observer's line of sight and the axis of the jet cone. We can describe that with an angle $\\theta_{\\rm obs}$ between both directions. It is very unlikely for an observer to be looking precisely at the direction of the jet axis (to have $\\theta_{\\rm obs}=0$). We get an 'on-axis' GRB and a 'regular' optical afterglow when $\\theta_{\\rm obs} < \\theta_{\\rm j}$. If the jet cone is directed so that $\\theta_{\\rm obs} > \\theta_{\\rm j}$,  an observer, due to the beaming effect, probably cannot observe the prompt gamma emission, i.e., detect a GRB. But the deceleration of the jet will gradually expand the beaming angle and the observer might detect the afterglow. The observed afterglow that was not preceded by a prompt gamma emission is called an orphan afterglow (OA)\\footnote{The OAs we are discussing here are not to be confused with optical afterglows, for which the observer does not detect the prompt gamma emission for some reason, even though the jet cone is directed into the observer's line of sight.}. Even though OAs have not been conclusively detected yet, they have been studied a great deal over the last 15 years (\\citealt{granot} and references therein).  As part of the scientific community preparation for the $\\textit{Gaia}$ mission, we discuss $\\textit{Gaia}$'s potential for detecting these short, rapidly fading and unpredictable transients. We present here results of the simulation we performed in order to estimate the number of possible GRB optical afterglow detections with the $\\textit{Gaia}$ satellite. The simulation is divided in two parts, the first one concerning on-axis GRB afterglows and the second one dealing with orphan afterglows. In Section 2 we briefly discuss the dynamics and basic characteristics of $\\textit{Gaia}$. The simulation with all the necessary parameters for both on-axis and orphan afterglows is described in Section 3. In Section 4, we present and discuss results of the simulation and possible ways of afterglow identification.  \\begin{figure}[!t] \\epsscale{1} \\plotone{f1.eps} \\caption{Sketch showing how $\\textit{Gaia}$ will scan the sky. The telescope axes are separated by an angle of $\\beta = 106.5^{\\circ}$. The satellite spins with a constant angular velocity $\\omega$ around its main axis, which is kept at an angle $\\gamma = 45^{\\circ}$ from the Sun at all times. The main axis experiences slow precessional motion around the Earth-to-Sun direction ($\\Omega$). Rotation around the Sun (angular velocity $\\rho$) is also taken into account. Details of the dynamics are given in the Appendix.} \\label{fig1} \\end{figure} ",
        "conclusions": "\\label{s5} We have shown that the number of expected GRB optical afterglow detections with $\\textit{Gaia}$ satellite during its five years of expected operation is of the order of a few tens. Most of the expected detections are of the orphan afterglows, while the possibility of an on-axis afterglow detection is small (less than 10). The results are dependent on a set of various parameters, especially in the case of OAs. Since we made some assumptions and simplifications in the course of our simulation, we would like to point out that these results have to be interpreted as an order-of-magnitude estimations. In addition to the small number of expected detections, poor time sampling of the afterglow light curve and problem of afterglow identification make the study of GRB optical afterglows with $\\textit{Gaia}$ difficult. On the other hand, since OAs have not been conclusively detected yet, a possible detection would provide valuable new data."
    },
    "1107/1107.5155_arXiv.txt": {
        "abstract": "{Narrow-line Seyfert 1 (NLS1) galaxies are a class of active galactic nuclei ( AGN) that have all the properties of type 1 Seyfert galaxies but show peculiar characteristics, including the narrowest Balmer lines, strongest Fe II emission, and extreme properties in the X-rays. Line and continuum radio observations provide an optimal tool to access the (often) optically obscured innermost regions of AGN and reveal the kinematics of the gas around their central engines.} {We investigate the interplay between the peculiar NLS1 class of AGN and the maser phenomenon, to help us understand the nature of the maser emission in some NLS1s where water maser emission has been detected.} {We observed a sample of NLS1 galaxies with the Green Bank Telescope in a search for water maser emission at 22 GHz. We also reduced and analysed archival Green Bank Telescope and Very Large Array data and produced 22-GHz spectra for the five NLS1 galaxies with detected maser emission. In particular, we imaged the maser and nuclear radio continuum of NGC~5506 at subarcsec scales with the Very Large Array.} {We discovered maser emission in two NLS1 galaxies: IGR~J16385-2057, and IRAS~03450+0055. In addition to the three previously known maser detections in the NLS1s Mrk~766, NGC~4051, and NGC~5506, this yields a water maser detection rate in NLS1 galaxies of $\\sim$ 7\\% (5/71). This value rises significantly to $\\sim$ 21\\% (5/24) when considering only NLS1 galaxies at recessional velocities less than 10000 \\kms. For NGC~4051 and NGC~5506, we find that the water maser emission is located within 5 and 12 pc, respectively, of nuclear radio continuum knots, which are interpreted as core-jet structures.} {The water maser detection rate in NLS1s is surprisingly high, much higher than the detection rate obtained for type 1 AGN and similar to those in Seyfert 2 and low-ionization nuclear emission-line region galaxies. The masers in NGC~4051 and NGC~5506 are nuclear and associated with the AGN, either with an accretion disk, a radio jet, or a nuclear outflow. The apparent lack of high-velocity maser features and evidence, recently reported, of radiative outflows and radio jets in the host galaxies seems to favour interpretation as a jet or an outflow. A similar association is also seemingly true for the maser in Mrk~766, IGR~J16385-2057, and IRAS~03450+0055, although, in these cases, without radio interferometric measurements we cannot rule out an off-nuclear origin of the emission.} ",
        "introduction": "According to the widely accepted unified model of active galactic nuclei (AGN; e.g., \\citealt{antonucci93}, \\citealt{urry95}), in their very centres there is a supermassive black hole surrounded by a parsec-scale accretion disk. The disk emits intensely in the ultraviolet (UV) and soft X-ray wavelengths. A torus (or a thick disk) of atomic and molecular gas, with a size of 1-100 pc, surrounds the accretion disk and obscures the optical and UV emission along certain directions. Therefore, depending on the line of sight, the object appears as either a type 1 or type 2 AGN. In type 1 AGN, the accretion disk and black hole are viewed through the hole in the torus, while in type 2 AGN they are obscured by the torus.  Narrow-line Seyfert 1 (NLS1) galaxies have attracted particular attention. While Seyfert 2's are active galaxies with narrow emission-line optical spectra, NLS1s have the broad emission-line optical spectra of type 1 Seyfert galaxies, but with the narrowest Balmer lines from the broad line region and the strongest Fe II emission (e.g., \\citealt{osterbrock89}, \\citealt{veron01}). Extreme properties are also observed in X-rays, such as a strong soft excess emission below 1 keV and rapid flux variability, accompanied by steeper photon indices with respect to (w.r.t.) Seyfert 1 (Sy\\,1) galaxies (e.g., \\citealt{leighly99}, \\citealt{gallo04}). The complex X-ray spectra often display signs of cold and ionized absorption, a partial covering of material along the line of sight, and/or (relativistic) reflection components from the inner side of the accretion disk (see \\citealt{komossa08} for a review). Because of its extreme and ambiguous characteristics, this class of AGN seems to challenge the Unified Model, hence deserves a more detailed investigation.  Fundamental studies of the central regions of AGN are complicated by the extremely small scales and complex structures of the nuclear components. Furthermore, in type 2 AGN in particular, the inner regions are often obscured at optical and UV wavelengths. Observations at infrared (IR, the band where most of the nuclear radiation absorbed by the torus is re-emitted), X-ray, and radio frequencies can, however, access these obscured regions. In particular, at radio wavelengths, water maser studies are a unique tool for investigating the structure and kinematics of the gas close to the black holes.  Most extragalactic \\water\\ masers are associated with AGN, where they have been related to three distinct phenomena. The first is that they may form directly in the nuclear accretion disk, where they can be used to trace the disk geometry and rotation velocities, and may reveal the enclosed nuclear mass (e.\\ g.\\ \\citealt{kuo11}). In some cases, they are also being used to measure accurate distances to their host galaxies (for the most recent case of UGC~3789, see \\citealt{braatz10}). Secondly, they are associated with radio jets, where they are either the result of an interaction between the jet(s) and a molecular cloud or a coincidental overlap along the line of sight between a warm, dense molecular cloud and the radio continuum of the jet (NGC~1052: e.\\ g.\\ \\citealt{sawada08}; NGC~1068: e.\\ g.\\ \\citealt{gallimore01}; Mrk~348: e.\\ g.\\ \\citealt{peck03}), hence provide important information about the evolution of jets and their hotspots. Thirdly, they are associated with nuclear outflows, tracing the velocity and geometry of nuclear winds, as in the case of Circinus \\citep{greenhill03}.  So far, more than 3000 galaxies have been searched for water maser emission and  detections have been obtained in about 150 of them (Braatz et al., in prep), the majority being Seyfert 2 (Sy\\,2) galaxies. In this paper, we report the results of a new Green Bank Telescope (GBT) survey to search for water maser emission in objects belonging to the NLS1 class of galaxies and, by analysing the main characteristics and line features of the 5 NLS1 galaxies where water maser emission has been detected in the present and past surveys, we discuss the relation between the maser phenomenon and NLS1 galaxies in the framework of the unified scheme of AGN. ",
        "conclusions": "\\label{disc}  \\subsection{The water maser detection rate in NLS1 galaxies} \\label{detection}  Identifying water maser emission associated with Sy\\,2 galaxies is not particularly surprising, given that the unified scheme for AGN requires an obscuring structure, which is probably an edge-on disk or torus, along the line of sight towards the nucleus of a type 2 AGN. This structure can provide a molecular reservoir and the amplification paths necessary for maser action. The relation between a type 1 Seyfert and the maser phenomenon is less obvious. According to the same paradigm, in type 1 objects, accretion disks and/or tori, if present at all, should be orientated face-on, making them less likely to produce detectable maser emission. Accordingly, out of the 150 water maser sources detected so far, maser emission has been detected in only very few type 1 Seyfert galaxies. Interestingly, two galaxies among these are classified as NLS1s: NGC~4051 and NGC~5506. In view of the above considerations, the association between water maser emission and the NLS1 classification is clearly intriguing. We searched for water maser emission with the GBT in a sample of 17 Seyfert galaxies (classified by \\citealt{masetti08}) discovered by INTErnational Gamma-Ray Astrophysics Laboratory (INTEGRAL) as sources of hard X-rays between 20 and 40 keV \\citep{bird10}. We detected water maser emission in one object (Fig.~1), IGR~J16385. Interestingly, this was the only NLS1 in the sample \\citep{masetti08}. Further details on a statistical study of this campaign will be presented elsewhere (Castangia et al. in prep).  Galaxy classifications are taken from \\citet{veron10}, with the exception of NGC~5506, which was (re)classified by \\citet{nagar02} using near-IR spectroscopy. Nagar et al. detected the permitted OI $\\lambda$1.1287 $\\mu$m line (with full width at half maximum $<$ 2000 \\kms) and the ``1 micron Fe II'' lines to identify the galaxy as a NLS1. Their finding raises the important issue that other galaxies may be misclassified, depending on the quality and type of spectral information available. In principle, according to \\citet{nagar02}, other X-ray bright and highly variable ``type 2'' Seyferts, such as NGC~5506, may actually be partially obscured NLS1s. This is just a cautionary note that needs to be kept in mind when dealing with studies similar to the one proposed here.  Including all past surveys (Braatz et al.\\ in prep) and our new detections in IGR~J16385 and IRAS~03450, a total of 71 NLS1 have been searched for water maser lines with five successful detections (Table~\\ref{all}). This yields a detection rate of $\\sim$ 7\\%, which is comparable to the global rates of AGN surveys (e.\\ g.\\ \\citealt{braatz97}). While this result is surprising, the maser detection rate in NLS1s becomes even more impressive when we consider a volume-limited sample that somewhat minimizes the limitation in sensitivity of the survey(s). When considering only NLS1 galaxies at recessional velocities less than 10000 \\kms, the detection rate goes up dramatically to $\\sim$ 21 \\% (5/24). This value approaches the highest detection rates ever obtained for similarly volume-limited samples, in any class of AGN. For example, the most prolific class of AGN, the Sy\\,2 galaxies, have detection rates of up to $\\sim$ 20\\%, while in `pure' Sy\\,1 galaxies the percentage of water maser detections never exceeds 1\\% (e.\\ g.\\ \\citealt{braatz04};\\cite{bennert09}; \\cite{zhang10}; \\citealt{ramolla11}).  \\subsection{The origin of water maser emission in NLS1 galaxies} \\label{origin}  \\begin{figure*} \\centering \\includegraphics[width= 17 cm]{17213fg3.eps} \\caption{1.6-GHz EVN radio continuum maps of the nuclear region of the NLS1 NGC~4051 (\\citealt{giroletti09}). The size of the beam is 13.5 $\\times$ 11.3 mas. The cross indicates the position (RA$_{J2000}$ = 12h03m09\\fs61; Dec.$_{J2000}$ = +44d31'52\\farcs7) of the maser derived with the VLA A-array (\\cite{hagi07}; present work). The circle represents the 0\\farcs1 arcsec VLA nominal position accuracy of the maser emission.} \\label{fig:n4051} \\end{figure*}  Water masers associated with an AGN accretion disk exhibit a typical spectrum with three distinct groups of features. One ``cluster'' of lines is found at the recessional velocity of the galaxy, while the other two are blueshifted and redshifted by hundreds of \\kms (e.g., for NGC~4258: \\citealt{miyoshi95}; for UGC~3789: \\citealt{braatz10}). Water masers associated with radio jets (like in M~51, NGC~1068, Mrk~348, and NGC~1052) instead show single, broad (with linewidths of $\\sim$ 100 km/s) maser lines offset from the systemic velocity of the host galaxies (e.g., see \\citealt{lo05} and references therein). In the Circinus galaxy, the only object where some of the maser emission has been confidently associated with a nuclear outflow, the maser features associated with the outflow have relatively narrower widths (a few \\kms), while line velocities are again found to be both red- and blue-shifted w.r.t the systemic velocity (\\citealt{greenhill03}).  The characteristics of the maser lines found in NLS1 galaxies (Sect.~\\ref{res}) and, in particular, the lack of clear evidence for high-velocity lines (Fig.~\\ref{fig:threespectra}) suggests that masers in NLS1s differ from the disk masers in classic Sy\\,2, perhaps being associated with jets and/or nuclear outflows. However, it has to be kept in mind that the maser lines associated with jets, especially in Mrk~348 and NGC~1052, are typically (much) broader than those in NLS1s (100 vs. 10 \\kms). In addition, the displacements of the line velocities from the systemic velocity in jet-maser galaxies are larger than those in NLS1s (120-150 vs 40-80 \\kms) and all towards the red, with the only exception of NGC~1068 (\\citealt{gallimore01}). Detailed studies of each case are thus necessary to unveil the origin of the maser emission in the five NLS1 targets of this work.  There are presently only two maser NLS1 galaxies with available interferometric measurements of both the radio continuum (with Very Long Baseline Interferometry; VLBI) and the maser emission (with the VLA): NGC~4051 and NGC~5506. For both galaxies, the nuclear radio continuum emission is resolved into compact knots that can be interpreted as core-jet structures (\\citealt{giroletti09}, for NGC~4051; \\citealt{middelberg04}, for NGC~5506).  In the following, we consider, in detail for these two galaxies, the association of the maser emission with either accretion disks, jets, or outflows and we investigate the main characteristics of the five maser spectra found in NLS1s and their host galaxies to more tightly constrain the nature of the maser emission. The investigation of these five cases constitute a first attempt to understand the possible relation between NLS1 objects and the maser phenomenon.  \\subsubsection{NGC~4051} The nuclear radio source detected at both 1.6 and 5 GHz with the European VLBI Network (EVN) by \\citet{giroletti09} is most easily described in terms of a jet-base/outflow phenomenon. The detection by Giroletti \\& Panessa of three aligned sub-millijansky components in NGC 4051 suggests an ejection process. In particular, the position of the EVN radio core is coincident with that derived for the \\water\\ maser by \\citet{hagi07} and ourselves, within 0\\farcs1 (Sect.~\\ref{res}), corresponding to 5 pc at the distance of NGC~4051 (Fig.~\\ref{fig:n4051}). This suggests that the water maser emission arises from a molecular disk or from a nuclear jet/wind associated with the brightest radio continuum knot. This latter option seems to be consistent with the existence of a weak, relatively extended ($\\sim$ 60 pc) jet in NGC~4051 (\\citealt{jones11}, \\citealt{king11}, and \\citealt{maitra11}) and a shocked outflow (\\citealt{pounds11} and \\citealt{vaughan11}). In addition, the single-dish maser spectrum of NGC~4051 is somewhat reminiscent of that in Circinus where some of the maser emission is associated with a nuclear outflow \\citep{greenhill03}.  {\\it {An alternative na\\\"{i}ve interpretation:}} While the jet/outflow association of the maser emission is indeed the most likely option, a speculative interpretation of the maser spectrum shown in Fig.~\\ref{fig:threespectra} can be made in terms of an accretion disk maser. The maser spectrum could be described by three main peaks at velocities of $\\sim$ 670, 720, and 770 \\kms (see Fig.~\\ref{fig:threespectra}) roughly centred on the systemic velocity. The system thus resembles the triple peak profile of accretion disk masers, albeit with much lower velocity offsets. The 'triple-peak profile' is also somewhat confirmed by the VLA A-array spectrum taken at a different epoch (Fig.~\\ref{fig:twospectra}). In this scenario, we would obtain a rotation velocity of $V_{\\rm rot} = V_{\\rm obs}\\,\\cdot\\,(\\sin^{-1} i)$ $\\sim$ 50$\\,\\cdot\\,(\\sin^{-1} i)$ \\kms. From \\cite{uttley04}, we obtain $M_{\\rm BH}$ = 5 $\\times$ 10$^{5}$\\,M$_{\\odot}$ as the mass  of the nuclear engine (Table~\\ref{lines}). Combining these two parameters and assuming Keplerian rotation (as in NGC\\,4258), this yields an angular size of 18\\,$\\cdot$\\,($\\sin^{-2} i$) mas (see \\citealt{tarchi07a} for details on a similar calculation), corresponding to a galactocentric radius of $R_{\\rm GC}$ $\\sim$ 0.9$\\,\\cdot\\,(\\sin^{2} i)$ pc. Assuming a nuclear accretion disk seen approximately edge-on ($i$ $\\sim$ 90$\\degr$), the radius of the accretion disk in NGC~4051 would be almost ten times larger than the one found, for example, in NGC\\,4258 (0.1 pc). Such a large accretion disk in a Seyfert galaxy has, so far, only been directly measured in the Seyfert 1.9 galaxy NGC~1194, which has an inner and outer radius for the disk of 0.54 and 1.3 pc, respectively \\citep{kuo11}. Alternatively, a possibly lower value for the inclination of the disk and/or for the black hole mass in NGC~4051 would result in a size of the disk more consistent with the one in NGC~4258. Follow-up VLBI observations have been requested aimed at shedding light on the aforementioned scenario.  \\begin{figure*} \\centering \\includegraphics[width= 17 cm]{17213fg4.eps} \\caption{{\\it Left panel}: VLA A-array K-band continuum image of the nuclear region of NGC~5506. Contour levels are -1, 1, 2, 3, 4, 5, 6, 7, 8, 9, 10, 11\\,$\\times$\\,2.1\\,mJy\\,beam$^{-1}$ (1$\\sigma$ rms = 0.7\\,mJy\\,beam$^{-1}$). The beam size is 0\\farcs12 $\\times$ 0\\farcs08. We show the positions of the optical center (blue circle, taken from \\citealt{adelman08}), water maser emission in our VLA DnC-array (green circle) and VLA A-array (red circle) maps, respectively. The size of the bars indicates the uncertainty in the positions derived. The dotted-dashed square indicates the area covered by the VLBI map of \\citet{middelberg04} and the magenta square the position of the brightest source, B0, in this map. {\\it Right panel}: the EVN and MERLIN 6 cm image of the nuclear region of NGC~5506 by \\citet{middelberg04}. Offsets are relative to RA$_{\\rm J2000}$ = 14$^{\\rm h}$ 13$^{\\rm m}$ 14\\fs87953, Dec$_{\\rm J2000}$ = --03$^{\\rm \\circ}$ 12$^{\\rm \\prime}$ 27\\farcs6799. The size of the beam, shown in the bottom-left corner, is 16.5 $\\times$ 8.0 mas.} \\label{fig:trybig} \\end{figure*}  \\subsubsection{NGC~5506:} Figure~\\ref{fig:trybig} shows the innermost region of NGC~5506 as seen from our VLA K-band image (left panel) and from the EVN+MERLIN\\footnote{Multi-Element Radio Linked Interferometer Network} C-band map of \\citet{middelberg04} (right panel). Overlaid on the K-band contour map are the positions of the main centres of activity in the galaxy, namely the optical nucleus and the brightest radio knot detected by  \\citet{middelberg04}, in addition to the location of the maser emission derived in our VLA DnC and A-array maps. Within the uncertainties, all positions are coincident. \\citet{middelberg04} proposed two main scenarios to describe their VLBI map: i) the strongest knot, B0, is a candidate location of the AGN core owing to its flat spectrum, compactness, and high brightness temperature ($\\sim$ 3.6 $\\times$ 10$^{8}$ K); ii) the AGN might lie between B0 and B1, hence these two bright components are the ``working surfaces'' where double-sided jets from the AGN impact the interstellar medium (ISM). The region of emission north of B0 and component B2 could be low-surface-brightness emission seen extending from the hot spots to a larger scale. The flat radio spectrum of B0 is indicative of optically thick synchrotron emission, while ``working surfaces'', as in compact symmetric objects (CSOs), have steeper spectra, which are typical of optically thin synchrotron emission, therefore favouring B0 as being the core. In any case, our spectral line observations indicate a position for the maser that is consistent with either one of the nuclear components, hence an association of the maser emission with an accretion disk, a radio jet, or a nuclear outflow. Indications of a one-sided jet and an outflow are found for this galaxy (\\citealt{xant10}, and references therein) and the blue-shifted velocities of the maser lines w.r.t. the systemic velocity may be more consistent with a jet or an outflow origin.  \\citet{guainazzi10} measured the inclination of the nuclear disk in NGC5506 to be $i$=40\\degr using the profile of the X-ray FeK$\\alpha$ emission line. This measurement seems also consistent with the value of the nuclear column density (N$_{\\rm H}$ $\\sim$ 3$\\pm$1\\,$\\times$\\,$10^{22}$ cm$^{-2}$) for this galaxy obtained by \\citet{bianchi03}, which is in-between those typically found for Sy\\,1 and Sy\\,2 galaxies (N$_{\\rm H}$ $\\le$ $10^{22}$ and N$_{\\rm H}$ $\\sim$ $10^{23}$ cm$^{-2}$, respectively; e.g. \\citealt{malizia09}). Such an inclination for the disk in NGC~5506 would be less favourable for producing detectable (disk) maser emission (see also Sect.~\\ref{detection}). Furthermore, if we also assume that such an inclination can still provide a long enough amplification path for maser action, we would obtain a probability (P($i$)=cos($i$)) of detecting water masers in AGN accretion disks of $\\sim$ 77\\%\\footnote{This estimate is based on simple geometrical considerations, thus assumes that all AGN accretion disks host maser emission and ignores the many physical parameters on which the detection of masers depends.}. This value is much higher than any detection rate obtained in AGN surveys. Under the same assumptions made before and in the (unlikely) case that in NLS1s we only have disk-masers, we can then derive the range of accretion disks inclinations allowed to account for the maser detection rate found in NLS1s of $\\sim$ 7\\% (or $\\sim$ 20\\% when considering only targets within 10000 \\kms; see Sect.~\\ref{detection}). With these input parameters, the inclination of the nuclear disk favourable to producing maser action in NLS1s should not be less than 86\\degr (or 79\\%). All this seems to support the conclusion that the maser in NGC~5506 is not associated with an accretion disk, favouring instead an association with the radio jet or a nuclear outflow, or an atypical maser case.  \\begin{figure} \\centering \\includegraphics[width= 9 cm]{17213fg5.eps} \\caption{VLA A-array X-band naturally-weighted continuum image of Mrk~766. Contour levels are -1, 1, 2, 4, 8, 16\\,$\\times$0.2\\,mJy\\,beam$^{-1}$ (1$\\sigma$ rms = 0.07\\,mJy\\,beam$^{-1}$). The beam size, shown in the lower-left corner is 0\\farcs28 $\\times$ 0\\farcs24.} \\label{n4253} \\end{figure}  \\subsubsection{Mrk~766} For Mrk~766, neither interferometric measurements of the maser emission nor a VLBI map of the nuclear radio continuum are available. From VLA A-array archival data, we have produced a radio continuum map at X-band of the nucleus of Mrk~766. The image shows a compact nuclear source of $\\sim$ 5 mJy (Fig.~\\ref{n4253}). We have been recently granted EVLA A-array time to observe the maser emission in Mrk~766. These measurements will allow us to compare the position of the maser with that of the nuclear radio continuum source(s), and thereby discriminate between a nuclear maser and a (less likely) maser associated with a star-forming region in the large-scale disk of the galaxy. We note an analysis of X-ray Multi-mirror Mission (XMM) time-resolved X-ray spectroscopic data, has found in Mrk~766 evidence of gas in a Keplerian orbit around a (super)massive black hole \\citep{turner06}.  \\subsubsection{IGR~J16385} For IGR~J16385, only the 1.4 GHz radio continuum map from the National Radio Astronomy Observatory VLA Sky Survey (NVSS) at 1.4 GHz is available. The NVSS map shows a radio source of $\\sim$ 10 mJy at the position of the optical peak of the galaxy. Expanded VLA (EVLA) B-array observations of the radio continuum (X-Band) and maser line (K-Band) in IGR~J16385 have been requested and await evaluation.  \\subsubsection{IRAS~0345} This galaxy has not been well studied. The nuclear radio continuum has been imaged by \\citet{thean00} at 8.4 GHz with the VLA A-array, showing a single, slightly resolved source with an integrated flux of 6.8 mJy (Thean et al. 2000; their Fig.~1). The galaxy was also part of the sample studied by \\citet{thompson09} using Spitzer mid-infrared (MIR) spectroscopic data and their conclusion was that its observational characteristics are consistent with a nucleus obscured by a compact and clumpy toroidal structure of dust. Owing to the lack of radio interferometric measurements, however, the association of the maser emission with the AGN or the nuclear starburst present in this galaxy cannot yet be determined. Indeed, IR spectroscopic data of a sample of Seyfert galaxies that includes IRAS~0345 were analysed by \\citet{imanishi04}. Their work indicates that nuclear starbursts occurring in the dusty tori of Seyfert galaxies may be physically connected to the central AGN. This view is quite different from that of the classical unification paradigm, in which the dusty tori 'simply' hide the central AGN of Sy\\,2 galaxies and reprocess AGN radiation as infrared dust emission in Seyfert galaxies.\\\\  While water masers are not typically detected in Sy\\,1 galaxies, our work demonstrates that they {\\it are} detected in NLS1 galaxies, at a rate comparable to Sy\\,2's.  Thus it is clear that the NLS1 class and the maser phenomenon are somehow associated.  The details of this association are unclear.  The next step to investigating the nature of masers in NLS1 galaxies would be to image the masers on parsec or sub-parsec scales with VLBI observations.   The low black hole masses measured in NLS1 galaxies have suggested that NLS1s have accretion rates close to the maximum allowed values, causing their observed extreme behaviour (\\citealt{marconi08}, and references therein). Owing to their low black hole masses, these objects may represent an important link between AGN and the elusive intermediate-mass black holes \\citep{komossa08}. Strong radiation-pressure outflows represents another possible explanation of the peculiar characteristics of NLS1 galaxies. In particular, estimates of black hole masses turn out to be severely underestimated when the effect of radiation pressure is neglected, thus making NLS1s more extreme than they actually are (e.g., \\citealt{marconi08}). Another possibility that may account for the NLS1 properties is that they may be viewed at angles intermediate between type 1 and type 2 Seyfert galaxies. Nuclear obscuration may thus play an important role in determining the observed properties of NLS1s.  Compared to the general AGN population, the five galaxies examined in our study have low black hole masses and high accretion rates\\footnote{These estimates, however, need to be interpreted with caution since they strongly depend on the method used and the precision to which the distance of the galaxy is known.}, consistent with their classification as NLSy1 galaxies (Table~\\ref{lines}). A direct association between the high detection rate of water masers in NLS1 and the two aforementioned peculiar quantities in this class of AGN cannot be presently assessed and requires a more complete approach. At this stage, it seems instead that there is a possible relation between the maser action and the presence of nuclear outflows detected in NLS1. As mentioned in Sect.~\\ref{origin}, the water maser in Circinus, which has similar spectral properties to some maser spectra in our NLS1s, is associated with an outflow produced by the AGN activity (\\citealt{greenhill03}). The association of the maser emission in NLS1 with favourably orientated outflows may be corroborated by the blue-shifted velocities of the main maser features w.r.t. the systemic velocity of the galaxy in four out of five of our targets (see Table~\\ref{lines} and Sect.~\\ref{origin}). As shown in Sect.~\\ref{origin} for the case of NGC~5506, it is instead more controversial, although not implausible, that the intermediate inclination of nuclear disks and tori proposed for NLS1s can be related to the production of water masers. However, the presence of nuclear optically thick gas and/or the partial covering of the nucleus, possibly related to the aforementioned disks/tori inclination, has been invoked to describe the X-ray characteristics of a number of NLS1s, including our target galaxies (\\citealt{guainazzi10} for NGC~5506, \\citealt{turner06} for Mrk~766, and \\citealt{terashima09} for NGC~4051). Given the relation between nuclear obscuration and maser emission reported for relatively large samples of maser galaxies (hence, mostly type 2 Seyfert) by, e.g., \\citet{greenhill08}, \\citet{zhang06a}, and \\cite{zhang10}, this could partly justify the relatively high maser detection rate in NLS1s, especially w.r.t. that in Sy\\,1."
    },
    "1107/1107.0911_arXiv.txt": {
        "abstract": "As part of our survey of galactic stellar halos, we investigate the structure and stellar populations of the northern outer part of the stellar halo in NGC~55, a member galaxy of the Sculptor Group, using deep and wide-field $V$- and $I$-band images taken with Subaru/Suprime-Cam. Based on  the analysis of the color-magnitude diagrams (CMDs) for red-giant-branch (RGB) stars, we derive a tip of RGB (TRGB)-based distance modulus to the galaxy of $(m-M)_0 = 26.58 \\pm 0.11 (d = 2.1 \\pm 0.1 {\\rm Mpc})$. From the stellar density maps, we detect the asymmetrically disturbed, thick disk structure and two metal-poor overdense substructures in the north region of NGC~55, which may correspond to merger remnants associated with hierarchical formation of NGC~55's halo. In addition, we identify a diffuse metal-poor halo extended out to at least $z \\sim 16$~kpc from the galactic plane. The surface-brightness profiles toward the $z$-direction perpendicular to the galactic plane suggest that the stellar density distribution in the northern outer part of NGC~55 is described by a locally isothermal disk at $z \\lsim 6$~kpc and a likely diffuse metal-poor halo with $V$-band surface brightness of $\\mu_{\\rm V} \\gsim 32$ mag arcsec$^{-2}$, where old RGB stars dominate. We derive the metallicity distributions (MDs) of these structures on the basis of the photometric comparison of RGB stars with the theoretical stellar evolutionary models. The MDs of the thick disk structures show the peak and mean metallicity of [Fe/H]$_{\\rm peak} \\sim -1.4$ and [Fe/H]$_{\\rm mean} \\sim -1.7$, respectively, while the outer substructures show more metal-poor features than the thick disk structure. Combined with the current results with our previous study for M31's halo, we discuss the possible difference in the formation process of stellar halos among different Hubble types. ",
        "introduction": "\\label{sec:intro}  Our understanding of how a stellar halo surrounding a disk galaxy like the Milky Way has formed is still enigmatic, because the observational information of this faint galactic component is yet limited, except for Local Group galaxies. Halo stars in the Milky Way are characterized by low metal abundance and high velocity dispersion and this extreme nature reflects the early chemo-dynamical evolution of the Milky Way. Extensive analyses of these stars have revealed, e.g., the dual nature of the halo structure (inner flattened halo in prograde rotation and outer spherical halo in retrograde rotation) \\citep[e.g.,][]{Carollo07}. Also, recent growing observational evidence suggests that the Milky Way halo has formed from an assembly process of many subsystems, as deduced from stream-like halo substructures \\citep[e.g.,][]{Yanny00}. Similar substructures have also been identified in M31's halo, including the Giant Stellar Stream \\citep{Ibata07}. Indeed, M31's halo provides a clear external view of ancient stellar populations: the halo is found to be extended more than 150~kpc \\citep{Raja05} and has characteristic spatial structure and metallicity distributions \\citep[e.g.,][]{Kalirai06,Ibata07,Tanaka10}. Another Local Group galaxy, M33, seems to have a rather smooth halo, although its disk orientation prevents the more refined picture of the halo. Thus, detailed studies of old stellar halos in various-type galaxies provide us with important clues to the understanding of galaxy formation.  Exploring stellar halos even beyond the Local Group is of great importance to accomplish our understanding of their generic nature and past formation history in different disk galaxies at different environments. Model investigations \\citep[e.g.,][]{Kauffmann96} predict that each disk galaxy has been developed through a different formation and evolutionary path: the collapse epoch, star formation history, and assembly rate of subsystems are not common and each stellar halo is expected to hold a different morphology, as other components (bulge and disk) differ along the Hubble sequence.  In this regard, an edge-on spiral SB(s)m galaxy, NGC~55, which is a member of the nearby Sculptor group that consists of approximately 30 galaxies \\citep{Cote97,Jerjen00}, is an excellent test-bed for this study: it provides an external perspective of a nearby late-type disk galaxy and yet is close enough to resolve individual stars. For example, NGC~55 shows asymmetric extra-planar morphology as derived from emission-line images \\citep{Ferguson96} and its nuclear source shows, based on far-infrared measurements, energetic star formation rate, which is of comparable luminosity to the bright star formation regions at the center of M33 and the Milky Way \\citep{Engelbracht04}. Based on the observational study of asymptotic giant branch (AGB) stars, \\citet{Davidge05} found that there are no evidence for a young or intermediate-age component in the extraplanar region such as 2~kpc off of the disk plane, suggesting that the population has ages of about 10~Gyr. On the other hand, \\citet{Tikhonov05} found an exponential spatial distribution of RGB stars between 2 and 7~kpc from the center along the minor axis, using HST/WFPC2. \\citet{Seth05a,Seth05b}, as part of a snapshot survey of 16 nearby edge-on late-type galaxies, observed the thin disk of NGC~55 using HST/ACS. They determined the fundamental properties of the resolved stellar population such as the distance, MD and spatial distributions of main-sequence, AGB and RGB stars in the inner part at $z\\lsim1.5$~kpc. However, since all the previous studies using HST are restricted to a small field-of-view, the global structure of an outer part of NGC~55 is yet unknown. Therefore, we conduct a high-resolution, deep wide-field survey of the halo of NGC~55 using Subaru/Suprime-Cam, to elucidate its fundamental properties for the first time.  The layout of this paper is as follows. In Section~\\ref{sec:data}, we present observations and summarize procedures for performing photometry for our Suprime-Cam data. Sections~\\ref{sec:dist}, \\ref{sec:map} and \\ref{sec:metal} are devoted to our results for quantitative analysis of the CMDs of NGC~55's stellar populations, including the distance estimate using the TRGB stars, the detection of substructures in NGC~55's halo from the stellar population maps, the radial profiles of the resolved stars, and their MDs using the theoretical isochrones. In Section~\\ref{sec:dis}, comparing with previous studies on stellar halos, we discuss the implications of our results and conclusions are drawn. ",
        "conclusions": "\\label{sec:dis}  Our Suprime-Cam observation of NGC~55 has revealed that this late-type galaxy holds extended galactic structures well beyond its bright disk component, namely TDE, TDW, Sub 1, Sub 2, and a diffuse halo. We discuss here the possible origin of these extended components in NGC~55.  As already shown in Section 5, the stellar population in each of these structures seems to be statistically different from each other, as deduced from the analysis of these MDs. However, taking into account the spatial distribution of Sub 1 and Sub 2 as well as the difference of stellar populations between TDE and TDW, there is a possibility that both Sub 1 and Sub 2 may be associated with some merging event which also gives rise to disk thickening in its east part. If this conjecture is the case, then the MD of TDE should be a combination of the MD of TDW and that of either Sub 1 or Sub 2, which is however unlikely as it follows from Figure 10. Thus, the difference in stellar populations between TDE and TDW may be attributed to another merging event, which is not related to the origin of Sub 1 and Sub 2.  Next, we compare the stellar populations of NGC~55's substructures with those of M31's halo. The MDs of NGC~55's substructures show a very different shape from those of M31's Giant Southern Stream, thereby suggesting that stellar population is quite different between NGC~55 and M31. Furthermore, the MDs of other substructures in M31 such as Stream~C and D with low stellar density \\citep{Tanaka10} are also different from those of NGC~55. These results may suggest the difference in the formation processes of stellar halos among different Hubble types of galaxies.  Previous studies of M31's stellar halo show a positive correlation between metal abundance and surface brightness for substructures observed in M31's halo \\citep{Gilbert09,Tanaka10}, i.e., tidal debris with higher surface brightness tend to be more metal-rich. This correlation implies that tidal debris with higher surface brightness may originate from more luminous and thus metal-rich dwarf satellites, and/or more recent encounters. For NGC~55's substructures, a similar correlation can be seen as shown in Figure~\\ref{fig:femu}, although the trend is systematically shifted toward a more metal-poor end than that for M31's halo. Thus, NGC~55's stellar halo may have formed through merging of relatively metal-poor dwarf galaxies compared to M31's stellar halo. Provided that the mass-metallicity relation of dwarf galaxies in the Sculptor Group is the same as that observed in the Local Group \\citet{Cote00}, NGC~55's halo may have originated from less massive dwarf galaxies than those for the formation of M31's halo, thereby implying that NGC~55 itself is less massive than M31 at the current epoch. This seems to be in agreement with the properties of M33 as well: this late-type galaxy has a less massive halo than M31 and has some faint substructures with metal-poor and low surface brightness in its halo \\citep{McConnachie10}.  The current study thus suggests that galaxies with different Hubble types have different properties of stellar halos as characterized by different metallicity and surface brightness in their substructures. However stellar halos in earlier Hubble types have yet been unidentified, simply because there are only few such early-type galaxies in the local volume so that current telescopes are able to resolve stars. On the other hand, a variety of substructures in the form of tidal streams have been found around more distant galaxies, where it is noteworthy that morphologies of tidal streams between early and late type galaxies are quite different \\citep[e.g.,][]{Peng02,Martinez10}. Resolving these substructures into individual stars by, e.g., Thirty Meter Telescope and/or Extremely Large Telescope, would provide further insights into the formation and evolution of stellar halos as a function of the Hubble sequence.   \\begin{figure}[htpd] \\begin{center} \\epsscale{1} \\plotone{./f11.eps} \\caption[a]{The median metallicity against surface brightness for the substructures of NGC~55 (filled black circles) and M31 (filled gray circles). Black crosses denote the disturbed and undisturbed thick disk structures of NGC~55. Metallicity uncertainties ($\\pm 0.2$ dex) are derived from the error analysis given in \\citet{Tanaka10}. It follows that tidal debris with higher surface brightness tend to be more metal-rich in both galaxies and that NGC~55's trend is systematically shifted towards a more metal-poor end than that for M31's halo. } \\label{fig:femu} \\end{center} \\end{figure}  \\medskip  As part of our survey of stellar halos in external galaxies, we have observed the north part of the stellar halo in NGC~55, using deep and wide-field $V$- and $I$-band images taken with Subaru/Suprime-Cam. On the basis of the analysis of the CMDs compared with theoretical isochrones, we have obtained the following major results:  \\begin{enumerate} \\item We have found that the stellar populations at $z \\gsim 5$~kpc above the disk plane are dominated by old RGB stars with lower metallicity than M31 and the Milky Way. We derive a TRGB-based distance modulus to the galaxy of $(m-M)_0 = 26.58 \\pm 0.11 (d = 2.1 \\pm 0.1 {\\rm Mpc})$.  \\item Based on the density map of NGC~55's stellar populations, we have found asymmetric thick disk features in the north region of NGC~55, which we refer to as TDE for a disturbed thick disk in the east and TDW for an undisturbed thick disk in the west. We have also identified two overdense substructures at $z \\sim 6.5$~kpc, which we refer to as Substructure~1 and 2 in this work. These substructures may correspond to remnants of merging events of small galaxies associated with the formation of NGC~55's halo.  \\item Detailed comparisons between the photometric data for halo regions of NGC~55 and that of the control field suggest the presence of a diffuse metal-poor halo extended out to at least $z \\sim 16$~kpc. However, it is yet unclear to what extent this faint component is actually distributed.  \\item The stellar density distribution perpendicular to the galactic plane of NGC~55 is described by a locally isothermal disk at $z \\lsim 6$~kpc and a diffuse metal-poor halo with $\\mu_{\\rm V} \\gsim 32$ mag arcsec$^{-2}$ at higher $z$, where old RGB stars dominate.  \\item From the photometric comparison of RGB stars with the theoretical stellar evolutionary model, we have obtained the MDs in the extraplanar regions. Both TDE and TDW show the peak, average and median metallicity of [Fe/H]$_{\\rm peak} \\sim -1.4$, [Fe/H]$_{\\rm mean} \\sim -1.7$ and [Fe/H]$_{\\rm med} \\sim -1.6$, respectively. In contrast, Substructure~1 and 2 show more metal-poor features than the thick disk structures. Furthermore, the low KS probabilities indicate that all the MDs are statistically different, suggesting that the stellar population in each substructure may have a different origin.  \\end{enumerate}"
    },
    "1107/1107.2914_arXiv.txt": {
        "abstract": "We report the results of a systematic survey for 95 GHz class I methanol masers towards a new sample of 192 massive young stellar object (MYSO) candidates associated with ongoing outflows (known as extended green objects or EGOs) identified from the \\emph{Spitzer} GLIMPSE survey. The observations were made with the Australia Telescope National Facility (ATNF) Mopra 22-m radio telescope and resulted in the detection of 105 new 95 GHz class I methanol masers. For 92 of the sources our observations provide the first identification of a class I maser transition associated with these objects (i.e. they are new class I methanol maser sources). Our survey proves that there is indeed a high detection rate (55\\%) of class I methanol masers towards EGOs. Comparison of the GLIMPSE point sources associated with EGOs with and without class I methanol maser detections shows they have similar mid-IR colors, with the majority meeting the color selection criteria -0.6$<$[5.8]-[8.0]$<$1.4 and 0.5$<$[3.6]-[4.5]$<$4.0. Investigations of the IRAC and MIPS 24 $\\mu$m colors and the associated millimeter dust clump properties (mass and density) of the EGOs for the sub-samples based on which class of methanol masers they are associated with suggests that the stellar mass range associated with class I methanol masers extends to lower masses than for class II methanol masers, or alternatively class I methanol masers may be associated with more than one evolutionary phase during the formation of a high-mass star. ",
        "introduction": "Methanol masers are quite common in massive star forming regions (MSFRs). Historically they have been empirically classified into two categories (class I and class II). The initial classification was based on the sources towards which the different transitions were detected (Batrla et al. 1988; Menten 1991a). Class II methanol masers are often found to be associated with strong centimeter continuum emission within 1$''$ (e.g. ultracompact (UC) regions; Walsh et al. 1998), infrared sources and OH masers. The strongest and best known class II transition is the 5$_{1}-6_{0}$ A$^{+}$ line at 6.7 GHz. Since its discovery (Menten 1991b), this class bright maser has become a reliable tool for detecting and studying regions where massive stars form and are in their very early stages of evolution (e.g. Minier et al 2003; Bourke et al 2005; Ellingsen 2006; Xu et al 2008; Pandian et al 2008). A number of surveys have been conducted for the 6.7 GHz class II methanol maser, resulting in the detection of $\\sim$ 900 sources in the Galaxy to date. The surveys include those that are summarized in the compilation of Pestalozzi et al. (2005) and the recent searches of Pandian et al. (2007), Ellingsen (2007), Xu et al. (2008; 2009) and the Parkes Methanol Multibeam (MMB) blind survey ($\\sim$ 300 sources (Green et al. 2009), but so far only part published by Caswell et al. 2010 and Green et al. 2010). In contrast class I methanol masers are usually found offset by 0.1 - 1.0 pc from UC H{\\sc ii} regions, infrared sources and OH masers (e.g. Plambeck \\& Menten 1990; Kurtz et al. 2004). A number of observations of class I methanol masers showed that they were located at the interface between molecular outflows and the parent cloud (Plambeck \\& Menten 1990; Johnston et al. 1997; Kurtz et al. 2004; Voronkov et al. 2006). These early results suggested that class I and class II methanol masers favored very different environments. These observational findings were supported by early theoretical models of methanol masers which found that class I masers are collisionally excited, in contrast to class II masers which have a radiative pumping mechanism (Cragg et al. 1992). Observations made since the mid-1990s (e.g. Slysh et al. 1994) have shown that at single dish resolutions class I and class II masers are often associated. Observations at high spatial resolution (e.g. Cyganowski et al. 2009) show that while the two types of masers are typically not co-spatial on arcsecond scales, they are usually driven by the same young stellar object.  Compared to class II masers, class I methanol masers are relatively poorly studied and understood. There have only been a few large surveys for class I masers (mainly at 44 and 95 GHz), primarily undertaken with single-dish telescopes (e.g. Haschick et al. 1990; Slysh et al. 1994; Val'tts et al. 2000; Ellingsen 2005) along with a few smaller scale interferometric searches (e.g. Kurtz et al 2004; Cyganowski et al. 2009). These have resulted in the detection of about 200 class I maser sources in the Galaxy to date (Val'tts et al. 2010).  Observations and theoretical calculations suggest that class I methanol masers can form in the interface region between outflows and the ambient molecular gas. The methanol abundance is significantly increased in the shocked interface regions (e.g. Gibb \\& Davis 1998) and the gas is heated and compressed providing more frequent collisions, which results in more efficient pumping (Voronkov et al. 2006). Interferometric observations have shown the location of some class~I methanol masers correlates closely with the shocked gas in outflows as traced by 2.12 $\\mu$m H$_{2}$ and SiO (e.g. Plambeck \\& Menten 1990 ; Voronkov et al. 2006). However, such a close physical association between the shocks driven by outflows and the class I masers has only been firmly established in a small number of sources. One of the problems encountered in studying the relationship between outflows (or shocks) and class I methanol masers has been in finding an appropriate outflow tracer (one which is frequently associated with class I maser emission). Recently Cyganowski et al. (2008) have suggested that the 4.5~$\\mu$m band of the \\emph{Spitzer Space Telescope's} Infrared Array Camera (IRAC) offers a promising new approach for identifying massive young stellar objects (MYSOs) with outflows. The strong, extended emission in this band is usually thought to be produced by shock-excited molecular H$_2$ and CO in protostellar outflows (e.g. Noriega-Crespo et al. 2004; Reach et al. 2006; Smith et al. 2006; Davis et al. 2007; Ybarra \\& Lada 2009, 2010). These extended 4.5~$\\mu$m emission features are commonly known as ``extended green objects'' (EGOs; Cyganowski et al. 2008) or ``green fuzzies'' (Chambers et al. 2009) due to the common color-coding of the 4.5~$\\mu$m band as green in three-color images. These objects are providing a new and powerful outflow tracer for MSFRs. Cyganowski et al. (2008) have catalogued over 300 EGOs from the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE) I survey (Churchwell et al. 2009), and they divided cataloged EGOs into ``likely'' and ``possible'' MYSO outflow candidates based primarily on the angular extent and morphology of the 4.5~$\\mu$m. Recent mid-infrared spectroscopic observations of two EGOs identified by Cyganowski et al. showed the presence of strong shocked H$_2$ in the 4-5 $\\mu$m wavelength range from one of the sources, but no evidence for shocked gas in the second (De Buizer \\& Vacca 2010). They suggest that some EGOs may be due to spatial variations in the mid-infrared extinction and exaggerated color stretches rather than shocked gas, however, the rate of mis-classification of EGOs is at present very poorly determined. Based on the mid-IR colors of EGOs and their correlation with infrared dark clouds (IRDC) and known 6.7 GHz methanol masers, Cyganowski et al. (2008) further suggests that most EGOs trace a population of actively accreting MYSO outflow candidates. Utilizing the published high resolution masers known at the time, Cyganowski et al. found that 6.7 GHz class II methanol masers are associated with 73\\% (35 detections from 48 EGOs) of ``likely'' and 27\\% (11 detections from 67 EGOs) of ``possible'' MYSO outflow candidate EGOs.  Chen et al. (2009) have analyzed 61 EGOs from the Cyganowski et al. catalog that have been included in class I methanol maser surveys. Four previous class I methanol maser surveys (for the 44 GHz $7_0$--$6_{1}$ A$^{+}$ transition by Slysh et al. (1994) and Kurtz et al. (2004) and for the 95 GHz $8_0$--$7_{1}$ A$^{+}$ transition by Val'tts et al. (2000) and Ellingsen (2005)) were included in their statistical analysis. They found that 41 EGOs are associated with one or both of the 95 and 44 GHz class I methanol masers within 1 arcmin, thus the expected detection rate of class I masers in EGOs is 67\\% at this resolution. Based on this high predicted detection rate, they suggested that GLIMPSE-identified EGOs might provide one of the best targeted samples for searching for class I methanol masers. However their statistical study is also subject to unknown influences produced by the target selection effects in the four class I maser surveys they have utilised in the study. For example, most of the currently known class I masers are associated with class II masers (72\\% as reported by Val\\'tts \\& Larinov 2007), for which a good association with EGOs has already been demonstrated by Cyganowski et al. (2008). Table 3 of Chen et al. (2009) shows that most EGOs (57/61) used in their statistical analysis are also associated with 6.7 GHz class II masers. According to Cyganowski et al. (2008), the majority of EGOs (35/48=73\\% in the likely sample) are associated with 6.7 GHz class II masers, and the sample of known class I masers is largely a subset of the sample of known class II masers, so a high detection rate for class I masers towards these EGOs is not unexpected. Recent 6.7 GHz class II and 44 GHz class I methanol maser surveys toward $\\sim$20 EGOs in the northern hemisphere with the VLA reported by Cyganowski et al. (2009) also show that $>$64\\% and 89\\% of EGOs are associated with class II and class I methanol masers, respectively. However, their observations only targeted a small number ($\\sim$ 20) of ``likely'' outflow candidate EGOs. And their 44 GHz survey sample is essentially a 6.7 GHz methanol maser selected sample (19 sources observed at 44 GHz (17/19 with class I masers), of which 18 had associated 6.7 GHz methanol masers (16/18 with class I masers)).  In order to determine whether there is indeed a high association rate of class I methanol masers in EGOs and to investigate the relationship between them, it is necessary to perform a class~I maser search towards a full EGO-based target sample. In this paper, we report a 95 GHz class I methanol maser survey towards an EGO-selected sample which has been undertaken with the Australia Telescope National Facility (ATNF) 22-m millimetre antenna at Mopra. In Section 2 we describe the sample and observations, in Section 3 we present the results of the survey, the discussion is given in Section 4, followed by our conclusions in Section 5. ",
        "conclusions": "Using the Mopra telescope, we have performed a systematic search for 95 GHz class I methanol masers toward EGOs. EGOs are new MYSO candidates with ongoing outflows identified from the \\emph{Spitzer} GLIMPSE I survey. We detected 105 new 95 GHz masers from a sample of 192 targets. Of these, 92 have no previously observed class I methanol maser activity, while the remaining 13 sources have been detected in the 44 GHz transition. Thus our single-dish survey proves that there is indeed a high detection rate ($\\sim$55\\%) of class I methanol masers in EGOs. Our findings increase the number of published class I methanol masers to 290 (an additional 92 on top of the $\\sim$198 from Val'tts \\& Larinov 2010). Mid-IR color analysis shows that the color-color region occupied by the GLIMPSE point sources for EGOs which are and are not associated with class I methanol masers are very similar, and mostly located in ranges -0.6$<$[5.8]-[8.0]$<$1.4 and 0.5$<$[3.6]-[4.5]$<$4.0 (see section 4.1 for detailed discussion of the uncertainties involved in this analysis). We find that the detection rate of class I methanol maser is likely to be higher in those sources with redder GLIMPSE point source colors.  Comparison of the [3.6]-[5.8] vs. [8.0]-[24] colors determined with integrated fluxes from Cyganowski et al. (2008) for the subsamples of the EGOs based on which class of methanol masers they are associated with, shows that those which are only associated class I methanol masers extend to less red colors than those associated with both classes of methanol maser. We suggest that the less red colors of class I methanol maser only EGOs is either because the class I only EGOs are associated with lower stellar mass objects, or because class I maser emission arises at more than one evolutionary phase of the high-mass star formation process. On the basis of current observations both scenarios can be plausibly argued and further observations will be required to determine which, if either of these hypotheses is correct. The thermal molecular line observations taken in conjunction with our maser search will be useful for trying to determine which of these scenarios is more likely. It will also be important to undertake high resolution observations of the class I maser emission in EGOs which are only associated with class I methanol masers to determine where the maser emission arises relative to the EGO. These observations are required to rule out the possibility that these sources represent a sample of chance associations between the EGOs and class I masers.  Analysis of the properties of mm dust clumps associated with class I methanol masers (for a subset of the EGOs in the class I maser survey sample which have available millimeter continuum data) shows that the luminosity of the class I methanol masers is correlated with the both the mass and density of the associated dust clump. The more massive and denser the clump, the stronger the class I methanol emission will be. We also find that the EGOs which are only associated with class I methanol masers have a lower maser luminosity and mass/density of dust clump. This finding supports either the hypothesis that the class I maser can trace a population with lower stellar masses, or that class I methanol masers may be associated with more than one evolutionary phase during the formation of a high-mass star."
    },
    "1107/1107.0316_arXiv.txt": {
        "abstract": "In this paper we present a model of the $\\beta$ Leo debris disc, with an emphasis on modelling the resolved PACS images obtained as part of the Herschel key programme DEBRIS. We also present new SPIRE images of the disc at 250 $\\micron$, as well as new constraints on the disc from SCUBA-2, mid-IR and scattered light imaging. Combining all available observational constraints, we find three possible models for the $\\beta$ Leo (HD102647) debris disc: (i) A 2 component model, comprised of a hot component at 2 AU  and a cold component from 15-70 AU. (ii) A 3 component model with hot dust at 2 AU, warm dust at 9 AU, and a cold component from 30-70 AU, is equally valid since the cold emission is not resolved within 30 AU. (iii)  A somewhat less likely possibility is that the system consists of a single very eccentric planetesimal population, with pericentres at 2 AU and apocentres at 65 AU. Thus, despite the wealth of observational constraints significant ambiguities remain; deep mid-IR and scattered light imaging of the dust distribution within 30 AU seems the most promising method to resolve the degeneracy. We discuss the implications for the possible planetary system architecture; e.g., the 2 component model suggests planets may exist at 2-15 AU, while the 3 component model suggests planets between 2-30 AU with a stable region containing the dust belt at 9 AU, and there should be no planets between 2-65 AU in the eccentric planetesimal model. We suggest that the hot dust may originate in the disintegration of comets scattered in from the cold disc, and examine all A stars known to harbour both hot and cold dust to consider the possibility that the ratio of hot and cold dust luminosities is indicative of the intervening planetary system architecture. ",
        "introduction": "Debris discs are distributions of dust and planetesimals with radii of 1-1000 AU around main sequence stars (See \\citealp{2008ARAA..46..339W} for a recent review). The dust grains in these discs are small, and so cannot be primordial as they would have been blown out of the system by radiation pressure on timescales shorter than the stellar age. The dust must be continually  replenished from a population of colliding planetesimals, thought to contain bodies up to $\\sim$ 1km in size (\\citealp{2002MNRAS.334..589W}). However, as debris discs age the planetesimal population is ground down, so discs become less massive and fainter (\\citealp{2003ApJ...598..626D}). At far-infrared and sub-millimetre wavelengths debris discs are optically thin, the disc to star contrast is favourable, and these wavelengths are sensitive to the large (up to $\\sim$1mm) grains that dominate the dust mass in debris discs.  The dust morphology of a debris disc can be shaped by planets in the system so resolved images of discs help constrain models of structure and evolution of planetary systems. Resolved images can indicate that infrared excess is being produced by multiple dust populations and can also break the degeneracy between the radial location of the dust and its temperature.   The DEBRIS (Disc Emission via a Bias-free Reconnaissance in the Infrared/Sub-mm) survey (Matthews et al. in Prep, \\citealp{ 2010MNRAS.403.1089P}), is an Open Time Key Program on the  \\emph{Herschel Space Observatory} which uses PACS (Photodetector Array Camera and Spectrometer \\citealp{2010AA...518L...2P}) and SPIRE (Spectral and Photometric Imaging REceiver, \\citealp{2010AA...518L...3G}) to detect, resolve and characterise debris discs around a volume-limited sample of 446 A through M stars (\\citealp{2010AA...518L.135M}). $\\beta$ Leo (HD 102647) was observed at 100 $\\micron$ and 160 $\\micron$ with PACS as part of the DEBRIS survey Science Demonstration Phase (\\citealp{2010AA...518L.135M}).  $\\beta$ Leo (A3V, $L_{*}$=14.0L$_{\\odot}$) is a $\\delta$ Scuti type star at a distance of 11.1 pc. The infrared excess around this main-sequence star was first discovered by \\emph{IRAS} (Infrared Astronomical Satellite \\citealp{1992AAS...96..625O}), then confirmed with \\emph{Spitzer} (\\citealp{2006ApJ...653..675S}). The excess was unresolved in mid-IR imaging (\\citealp{2001AJ....122.2047J,2009ApJ...691.1896A}, $\\S$ 2.2) although differences between the \\emph{IRAS}  and \\emph{ISO} (Infrared Space Observatory) fluxes led \\citet{2002AA...387..285L} to suggest that the disc emission may be somewhat extended in the ISO beam (52'' aperture). \\citet{2006ApJS..166..351C} obtained a Spitzer IRS (Spitzer Infrared Spectrograph) spectrum of $\\beta$ Leo (See $\\S$ 3.2) and found a featureless continuum spectrum consistent with dust at $\\sim$120 K located at 19 AU from the star. A very hot excess has also been partially resolved using infrared interferometry with the FLUOR (Fiber Linked Unit for Recombination) instrument at the CHARA (Center for High Angular Resolution Astronomy) array at 2 $\\micron$ and BLINC (Bracewell Infrared Nulling Cryostat) at 10 $\\micron$ (\\citealp{2009ApJ...691.1896A}, Stock et al. (2010), see $\\S$ 2.8 for more details).  $\\beta$ Leo is thought to be a member of the IC2391 moving group, giving an age of 45 Myr (\\citealp{2010AJ....140..713N}). The age of this source determined from isochrone fitting  is  50-331 Myr  (\\citealp{1999AA...348..897L,2001ApJ...546..352S}). \\citet{2004AA...426..601D} derive an age from measuring the stellar radius and suggest 100 Myr. We assume an age of 45 Myr for this paper. $\\beta$ Leo does not have any known companions within a few arcsecs of the star. The Washington Double Star (WDS) catalogue lists three companions with common proper motion for $\\beta$ Leo. These stars are located from 40$\\arcsec$ to 240$\\arcsec$ from the primary with V magnitude differences of 6.3 to 13 (\\citealp{1997A&AS..125..523W}); these stars are not, however, physically associated with $\\beta$ Leo (\\citealp{2010MNRAS.403.1089P}). $\\beta$ Leo was also included in a high precision radial velocity survey of early type dwarfs with HARPS (\\citealp{2009AA...495..335L}) for planetary or brown dwarf mass candidates and was found to have no companions with mass >4.2M$_{jup}$ with periods <10 days with 99.7 percent probability.  This paper uses multi-wavelength modelling of the $\\beta$ Leo debris disc including PACS and SPIRE observations to present a self-consistent explanation of the system. In \\S2 we present the observations, including 100 and 160 $\\micron$ PACS imaging (previously published in \\citealp{2010AA...518L.135M}) and archive Gemini MICHELLE 12 $\\micron$ and 18 $\\micron$ imaging. In \\S3 we confront the observations with models to determine the disc parameters. In \\S4 we discuss the implication of the inferred structure for the status of planet formation in this system. ",
        "conclusions": "We have presented detailed modelling of the debris disc around the 45 Myr old A star $\\beta$ Leo. We considered multi-wavelength data  to construct a complete picture of this source. Resolved images taken at 100 and 160 $\\micron$ using PACS on Herschel as part of the DEBRIS survey were used to place observational constraints on the radial location of the cold dust in this system. Resolving the cold dust is key to breaking the degeneracies inherent in SED modelling. We also use detection limits from unresolved images  at 12 and 18 $\\micron$ from MICHELLE on Gemini, at 0.6 $\\micron$ with ACS on HST, the unresolved Herschel SPIRE image at 250 $\\micron$  and detailed SED modelling including all data from the literature to gain a complete picture of the disc. Modelling indicates that for a 2 component model of the system consisting of a hot and cold disc, the cold disc imaged with Herschel PACS lies between 15-70 AU, with a surface density profile $\\Sigma \\propto r^{-1.5}$ at an inclination of 55$^{\\circ}$ from edge on, with the hot dust at 2 AU. SED fitting to observations from 5 $\\micron$ to 1mm, including the IRS spectrum (\\citealp{2006ApJS..166..351C})  suggest that the grain size distribution is consistent with that for a theoretical collisional cascade with a fixed maximum grain size of 1cm, a minimum grain size of 3 $\\micron$ (0.3$\\times D_{bl}$  for a composition with a silicate fraction of 20\\%, a porosity of 20 \\% with the rest of the grain composed of organic refractories. For the hot component the minimum grain size is 0.6 $\\micron$ (0.1$\\times D_{bl}$) and the silicate fraction was 60\\%.  A 3 component model indicates that another possible configuration consist of an inner edge to the cold disc at 30 AU, a warm ring at 9 AU and hot dust at 2 AU. It is also possible to fit the observations of $\\beta$ Leo using a single, very eccentric (e=0.92) planetesimal population, after \\citet{2010MNRAS.402..657W}. This degeneracy implies that even a wealth of multi-wavelength data including the resolved images may not be enough to uniquely constrain the location of the dust when there are multiple populations and the edges of the belts have not been resolved.  $\\beta$ Leo is a similar age to Fomalhaut and Vega, which also have hot dust within a few AU of the star. However, its cold disc is smaller and closer to the hot emission that either of these systems, and there may be an intermediate warm component. The most analogous disc among those resolved around A stars is Eta Tel, a 12 Myr star with a resolved disc from 21-26 AU and hot, unresolved dust at <4 AU (\\citealp{2009AA...493..299S}), but in terms of the separation and ratio of the hot and cold components as shown in Figure \\ref{A_stars}, $\\beta$ Leo most resembles HD71155.  We have also examined the population of A stars known to have both hot and cold dust discs, and we see two main 'types' of distributions between the hot and cold populations. First, discs such as  Vega and $\\beta$ Leo which have a large (>20 AU) separation between the location of the hot and cold dust, and whose hot dust has a high fractional luminosity. The other type of disc, typified by Eta Tel and HR8799 show an almost flat or, for HR8799, declining line in Figure \\ref{A_stars}, indicating that the hot dust has a similar or lower fractional luminosity than the cold component. In the case of HR8799 the hot and cold components are separated by a known planetary system  (\\citealp{2008Sci...322.1348M}) which may inhibit the transfer of material from the outer to inner regions of the system.  However, the small number of stars in this sample and the possibility of unknown systematic issues from comparing dust detected through different methods limit the conclusions that can currently be made."
    },
    "1107/1107.4645_arXiv.txt": {
        "abstract": "The existence of a bullet cluster (such as 1E0657-56) poses a challenge to the concordance $\\Lambda$ cold dark matter model.  Here we investigate the velocity distribution of dark matter halo pairs in large $N$-body simulations with differing box sizes ($250\\himpc - 2h^{-1}$Gpc) and resolutions.  We examine various basic statistics such as the halo masses, pairwise halo velocities ($v_{12}$), collisional angles, and pair separation distances. We then compare our results to the initial conditions required to reproduce the observational properties of 1E0657-56 in non-cosmological hydrodynamical simulations.  We find that the high velocity tail of the $v_{12}$ distribution extends to greater velocities as we increase the simulation box size. We also find that the number of high-$v_{12}$ pairs increases as we increase the particle count and resolution with a fixed box size, however, this increase is mostly due to lower mass halos which do not match the observed masses of 1E0657-56.  We find that the redshift evolution effect is not very strong for the $v_{12}$ distribution function between $z$=0.0 and $z$$\\sim$0.5.  We identify some pairs whose $v_{12}$ resemble the required initial conditions, however, even the best candidates have either wrong halo mass ratios, or too large separations. Our simulations suggest that it is very difficult to produce such initial conditions at $z=0.0, 0.296,$ \\& 0.489 in comoving volumes as large as (2\\,$h^{-1}$Gpc)$^{3}$. Based on the extrapolation of our cumulative $v_{12}$ function, we find that one needs a simulation with a comoving box size of (4.48\\,$h^{-1}$\\,Gpc)$^3$ and $2240^3$ DM particles in order to produce at least one pair of halos that resembles the required $v_{12}$ and observed masses of 1E0657-56. From our simulated $v_{12}$ probability distribution function, we find that the probability of finding a halo pair with $v_{12}\\geq 3000$\\,km\\,s$^{-1}$ and masses $\\geq 10^{14} \\Msun$ to be $2.76\\times10^{-8}$ at $z$=0.489. We conclude that either 1E0657-56 is incompatible with the concordance $\\Lam$CDM universe, or the initial conditions suggested by the non-cosmological simulations must be revised to give a lower value of $v_{12}$. ",
        "introduction": "\\label{sec:intro}  It is widely believed that the structure formation in our Universe is largely driven by the gravity of dark matter. Therefore it is worthwhile to probe dark matter dynamics through measurements of galaxy peculiar velocities and constrain our cosmological model by comparing against numerical simulations. In fact there has been extensive work along these lines, recovering the local density field from the measured velocity field \\citep{Bert89,Davis96,Willick96}. Unfortunately the observations of peculiar velocity fields contain large uncertainties, and accurate determination of the cosmological mass density parameter $\\Om$ turned out to be difficult using this method.  More recently, clusters of galaxies have been used to prove the existence of dark matter itself, thanks to accurate measurements of projected dark matter density using weak and strong lensing techniques. Some clusters show signs of a cluster-cluster merger, where the baryonic component and  collisionless dark matter show different spatial distributions, strongly supporting the existence of dark matter. Furthermore, using the shock features seen in the gas, one can infer the collision velocity of two galaxy clusters \\citep{Clowe04,Clowe06,Bradac06}. These new observations have brought renewed interest to dark matter dynamics and using it to check the standard $\\Lam$ cold dark matter ($\\Lam$CDM) cosmological model \\citep{Efsta90,Ostriker95}.  In particular, the observations of the massive cluster of galaxies 1E0657-56 seem to suggest a much higher relative dark matter halo velocity than one would expect in the $\\Lambda$CDM model.  This system includes a massive sub-cluster (the \"bullet\") with $M_{\\rm{bullet}} \\simeq 1.5 \\times 10^{14} \\Msun$ that has fallen through the parent cluster of $M_{\\rm{parent}} \\simeq 1.5 \\times 10^{15} \\Msun$ roughly $150$ million years ago, and is separated by $\\simeq 0.72$\\,Mpc on the sky at an observed redshift of $z$=0.296 \\citep{Clowe04,Clowe06,Bradac06}. The uniqueness of this system comes from the collision trajectory being almost perpendicular to our line of sight.  This provides an opportunity to better study the dynamics of large cluster collisions.  The Chandra observations revealed that the primary baryonic component had been stripped away in the collision and resided between the two clusters in the form of hot X-ray emitting gas \\citep{Mark06}.  This provides strong evidence for the existence of dark matter (DM); As the two clusters passed through each other, the baryonic components interacted and slowed down due to ram pressure, while the dark matter component was allowed to move ahead of the gas since it only interacts through gravity without dissipation. One can infer the velocity of the bow shock preceding the `bullet' using the shock Mach number and a measurement of the pre-shock temperature.  The inferred shock velocity was found to be $v_{\\rm{shock}} =4740^{+710}_{-550}$\\,km\\,s$^{-1}$ \\citep{Mark06}.  \\citet{Hayashi06} examined the Millennium Run \\citep{Mill05} in search for such a sub-cluster moving with a velocity relative to its parent cluster of $v_{\\rm{bullet}}=4500^{+1100}_{-800}$ km s$^{-1}$ \\citep{Mark04}.  Due to the limited volume of the simulation (500 $h^{-1}$Mpc)$^{3}$, few halos had masses comparable to 1E0657-56.  Still they estimated that about 1 in 100 have velocities comparable to the bullet cluster, and concluded that the event is not impossible within the current $\\Lambda$CDM model.  It is often assumed that the inferred shock velocity is equal to the velocity of the dark matter `bullet' itself.  Several groups have shown, however, that this is not necessarily true through the use of non-cosmological hydrodynamic simulations. \\citet{Milo07} used two dimensional simulations to find that the subcluster's velocity differed from the shock velocity by about 16\\%, bringing the relative velocity of DM halos down to $\\sim$\\,3980\\,km\\,s$^{-1}$.  They assumed a zero relative velocity at a separation distance of 4.6\\,Mpc for their initial conditions.  They also emphasized that their conclusion is sensitive to the initial mass and gas density profile of the two clusters. \\citet{Springel07} was able to reproduce the inferred shock velocity through the use of an idealized three dimensional hydrodynamic simulation with initial conditions that assumed a relative velocity of 2057\\,km\\,s$^{-1}$ at a separation distance of 3.37\\,Mpc, and found that the subcluster was moving with a relative speed of only $\\sim$\\,2600\\,km\\,s$^{-1}$ just after the collision. \\citet{Mast08} argued that \\citet{Springel07} failed to reproduce the observed displacement of X-ray peaks that represent an important indicator of the nature of the interaction.  In their simulations they found that in order to reproduce the observational data of 1E0657-56 a relative halo infall velocity of $\\sim$\\,3000\\,km\\,s$^{-1}$ at an initial separation distance of 5\\,Mpc was required.  Similar to previous work by \\citet{Hayashi06}, \\citet{Lee10} quantified the likelihood of finding bullet-like systems in the large cosmological N-body simulation MICE \\citep{MICE10}.  They examined DM halos at $z$=0.0 \\& 0.5, searching for a halo pair matching the initial conditions of \\citet{Mast08}.  They concluded that $\\Lambda$CDM is excluded by more than 99.91\\% confidence level at $z$=0.  Their results at $z$=0.5 are inconclusive due to limited statistics.  However, by fitting their pairwise velocity probability distribution function to a Gaussian distribution, they were able to estimate the probability of finding a pair with $v_{12}>3000$ km s$^{-1}$ to be $3.6\\times10^{-9}$ and $v_{12}>2000$ km s$^{-1}$ to be $2.2\\times10^{-3}$ at $z$=0.5.  They did warn that one must be careful about this approach since they are probing the tail of the distribution where their fits may not be accurate.  Most recently \\citet{Forero10} approached the problem from a different perspective. They studied data from the MareNostrum Universe \\citep{MareNostrum} which contains baryonic matter in addition to collisionless DM. Instead of examining the pairwise velocities of DM halo pairs, they concerned themselves with the physical separation between the dominant gas clump and its predominant DM structure. They argued that their approach provides a more robust comparison to observation; deriving the relative velocity from the observations includes statistical and systematic uncertainties whereas the separation uncertainty is dominated by statistical errors in the measuring process. Additionally they point out that current simulations do not include the proper prescriptions for cooling, star formation, or feedback which implies that their predictions of the detailed X-ray properties of hot gas in massive halos are not robust. Using their method they found that large displacements between gas \\& DM are common in $\\Lambda$CDM simulations therefore, 1E0657-56 should not be considered a challenge.  In this paper, we take a similar approach to that of \\citet{Lee10}, and examine large $\\Lambda$CDM $N$-body simulations to see how common these high relative velocities are among massive DM halos. One of the things that the earlier works have not performed is an examination of resolution and box size effect. Therefore we first conduct a study to determine the effects of increasing resolution or varying box sizes on the parameters of interest. We then examine our largest simulation in search for a pair matching the initial conditions required by \\citet{Mast08} to reproduce the observed properties of 1E0657-56.  The rest of the paper is organized as follows: Section\\,\\ref{sec:simulation} discusses simulation parameters, Section\\,\\ref{sec:analysis} shows the simulation results and examines the distribution of parameters relevant to this study, such as halo masses, pairwise velocity, and pair separation distances. Section\\,\\ref{sec:earlyz} examines the simulation results at earlier redshifts of $z$=0.296 \\& $z$=0.489, and how they relate to the bullet system. Finally, Section\\,\\ref{sec:discussion} contains concluding remarks and discussion of future prospects. ",
        "conclusions": "\\label{sec:discussion}  We performed many $N$-Body cosmological simulations with varying box sizes and resolutions in order to examine how changing these parameters affect the search for high-$v_{12}$ halo pairs comparable to the initial conditions required to reproduce the observed properties of the 1E0657-56 system in non-cosmological simulations. Using our largest L2016N1008 run, we examined the pairwise velocities, halo masses, and halo separation distances at $z$=0.0, 0.296, \\& 0.489.  We find that the high-$v_{12}$ tail of the distribution extends to a greater velocities as we increase the simulation box size. We also find that the number of high-$v_{12}$ pairs increased as we increase the particle count and resolution with a fixed box size, however, this increase is mostly due to lower mass halos which do not correspond to the characteristics of 1E0657-56.  We find that the redshift evolution effect is not very strong for the $v_{12}$ distribution function.  As we show in Table~\\ref{table:v12pairs}, some of the halo pairs have a high relative velocity similar to the initial conditions required to reproduce the observational quantities of 1E0657-56 in non-cosmological simulations, but they are galaxy group-scale halos (10$^{13}$$-$10$^{14}\\Msun$) and much less massive than the observed estimates for 1E0657-56.   We find that, in $N$-body simulations with comoving volumes of less than (2\\,$h^{-1}$\\,Gpc)$^3$, it is very difficult to reproduce a system that resembles the initial conditions required to reproduce the observational properties of 1E0657-56.  Based on the extrapolation of our cumulative $v_{12}$ function, we find that one needs a simulation with a comoving box size of (4.48\\,$h^{-1}$\\,Gpc)$^3$ and $2240^3$ DM particles in order to produce at least one pair of halos that resembles the initial conditions suggested by \\citet{Mast08}. In the future it would be useful to run larger simulations (e.g., with $\\sim$5\\,Gpc box and $\\sim$2500$^3$ particles) to improve the statistics of massive halos.  From the simulated $v_{12}$ PDF of halos, we calculated the probability of finding a halo pair with $v_{12}\\geq 3000$\\,km\\,s$^{-1}$ and masses $\\geq 10^{14} \\Msun$ to be $2.76\\times10^{-8}$, which is somewhat larger than previous work by \\citet{Lee10}.  However, both probabilities are quite small and the difference is negligible. These results suggest that a system like 1E0657-56 is currently incompatible with the concordance $\\Lam$CDM universe, if its initial condition really requires an initial pairwise velocity of $v_{12}\\geq 3000$\\,km\\,s$^{-1}$. As \\citet{Lee10} discussed in detail, there seems to be more systems like 1E0657-56 being observed already, which exacerbates the incompatibility in terms of probability. One other possibility is that there is something wrong with the referred non-cosmological simulations, and the suggested initial $v_{12}$ must be revised to a lower value.    \\begin{figure} \\includegraphics[scale=0.43]{fig15.eps} \\caption{Pairwise velocity probability distribution function for halo pairs with masses above 10$^{14} \\Msun$ in our L2016N1008 run.  The blue circles represent $v_{12}$ binned PDF data, the blue curve is the linearly interpolated values, and the red curve is the best-fit skew normal distribution \\citep{SkewNormal}.  Integrating the fit from $v_{12}=3000$\\,km\\,s$^{-1}$ to infinity gives $P(>3000$\\,km\\,s$^{-1})=2.8\\times10^{-8}$.  This very low probability suggests that it is very difficult to produce a halo pair with high mass and high-$v_{12}$ as the observed 1E0657-56. } \\label{fig15} \\end{figure}"
    },
    "1107/1107.3586_arXiv.txt": {
        "abstract": "{ In a previous study, we found that the detached post-common-envelope binary LTT 560 displays an H$\\alpha$ emission line consisting of two anti-phased components. While one of them was clearly caused by stellar activity from the secondary late-type main-sequence star, our analysis indicated that the white dwarf primary star is potentially the origin of the second component. However, the low resolution of the data means that our interpretation remains ambiguous. We here use time-series UVES data to compare the radial velocities of the H$\\alpha$ emission components to those of metal absorption lines from the primary and secondary stars. We find that the weaker component most certainly originates in the white dwarf and is probably caused by accretion. An abundance analysis of the white dwarf spectrum yields accretion rates that are consistent with mass loss from the secondary due to a stellar wind. The second and stronger H$\\alpha$ component is attributed to stellar activity on the secondary star. An active secondary is likely to be present because of the occurrence of a flare in our time-resolved spectroscopy. Furthermore, Roche tomography indicates that a significant area of the secondary star on its leading side and close to the first Lagrange point is covered by star spots. Finally, we derive the parameters for the system and place it in an evolutionary context. We find that the white dwarf is a very slow rotator, suggesting that it has had an angular-momentum evolution similar to that of field white dwarfs. We predict that LTT 560 will begin mass transfer via Roche-lobe overflow in $\\sim$3.5 Gyrs, and conclude that the system is representative of the progenitors of the current population of cataclysmic variables. It will most likely evolve to become an SU UMa type dwarf nova. } ",
        "introduction": "\\defcitealias{tappertetal07-1}{Paper I}% Cataclysmic variables (CVs) are thought to form from detached white dwarf (WD) / main-sequence star (MS) binaries that have experienced a common-envelope (CE) phase \\citep[e.g., ][]{taam+ricker10-1,webbink08-1}. In these ``post-common-envelope binaries\" (PCEBs), the white dwarf represents the more massive component in the system and is therefore called the primary, while the usually late-type (K--M) main-sequence star is known as the secondary. After the end of the CE phase, angular-momentum loss due to magnetic braking and/or gravitational radiation continues to decrease the separation between the two stars. This eventually brings the secondary's Roche lobe into contact with the stellar surface, thus initiating stable mass-transfer via Roche-lobe overflow and the semi-detached CV phase \\citep[][ and references therein]{ritter08-1}.  There is growing evidence that the accretion of material from the secondary star does not start with Roche-lobe overflow. The discovery of so-called ``low accretion rate polars'' \\citep{schwopeetal02-1}, which are likely progenitors of magnetic CVs \\citep{schmidtetal05-1,schmidtetal07-1}, and the detection of metal absorption lines in the UV spectra of non-magnetic PCEBs \\citep[e.g., ][]{kawkaetal08-1}, indicate that wind accretion from the active chromosphere of the secondary star is a common phenomenon in detached, short-period, WD/MS binaries.  \\citet[][ hereafter Paper I]{tappertetal07-1} found the H$\\alpha$ emission line in the detached PCEB \\object{LTT 560} to be a combination of two components. The radial velocity variations in the stronger H$\\alpha$ component agree well with those shown by the TiO absorption, and was thus identified as originating in the secondary star. The radial velocities of the weaker component showed an -- within the errors -- anti-phased behaviour with respect to the secondary star exhibiting a lower amplitude, and therefore had to be produced on the side of the centre-of-mass opposite to the secondary. However, the low spectral resolution of the data, and the contamination of the white-dwarf Balmer absorption lines with emission cores from the secondary, impeded velocity measurements of other, unambiguously intrinsic, spectral features from the white dwarf. Thus, the true origin of the weak emission component remained unresolved. The low temperature of the white dwarf in LTT 560 ($T \\sim 7500~\\mathrm{K}$, Paper I) and the evidence of ongoing accretion imply that there is a high probability that there are narrow metal absorption lines in the white-dwarf spectrum \\citep{zuckermanetal03-1}. The radial velocity curve of such lines should be undisturbed and thus faithfully track the motion of the white dwarf, motivating the high-resolution study presented in this paper. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[angle=270,width=1.0\\columnwidth]{16436f10.eps} \\caption[]{\\label{f-vsini} A close-up of three Fe\\,I absorption lines in the white dwarf photosphere. Overplotted on the average spectrum are three TLUSTY/SYNSPEC models for $T_\\mathrm{WD}=7500$\\,K, $\\log g=7.75$, and the abundances listed in Table\\,\\ref{t-abundances}. The solid line shows a non-rotating white dwarf, the dotted line corresponds to $v\\sin i=15$\\,km/s, and the dashed line to $v\\sin i=30$\\,km/s. } \\end{figure}  We have used Echelle spectroscopy of the PCEB LTT 560 to derive precise system parameters. From the metal abundances in the white dwarf photosphere, we determined an accretion rate of $\\sim5\\times10^{-15}\\Msun\\mathrm{yr^{-1}}$. \\citet{debes06-1} carried out a similar study of six white dwarf and main sequence binaries, both short-period (post-common envelope binaries, PCEBs) and wide binaries, and found accretion rates in the range $2\\times10^{-18}$ to $6\\times10^{-16}\\Msun\\mathrm{yr^{-1}}$. Their study includes the system RR\\,Cae, a white dwarf plus M-dwarf binary that resembles LTT\\,560 in many aspects. Both systems have nearly identical white dwarf masses and similar companion star masses, but the orbital period of RR\\,Cae is nearly double that of LTT\\,560. \\citet{debes06-1} determined an accretion rate of $4\\times10^{-16}\\Msun\\mathrm{yr^{-1}}$ for RR\\,Cae, and inferred, assuming a spherically symmetric wind and Bondi-Hoyle type accretion, a wind mass-loss rate of $\\dot M_\\mathrm{dM}=6\\times10^{-15}\\,\\mathrm{M_\\odot\\,yr^{-1}}$ for the companion star. Taking the Bondi-Hoyle scenario at face value, $\\dot M_\\mathrm{acc}\\propto1/R^2$, where $R$ is taken as the distance between the white dwarf and the M dwarf. For an equally large wind mass-loss rate, the accretion rate in LTT\\,560 would be about three times higher than in RR\\,Cae, which is only a factor of about three below the accretion rate that we estimated from the photospheric abundances. Considering all the uncertainties involved in estimating the accretion rate and mass-loss rate, RR\\,Cae and LTT\\,560 appear to behave rather similar, and we speculate that it is the higher accretion rate in LTT\\,560 that is responsible for the chromosphere/corona around the white dwarf that we observe in H$\\alpha$.  The accretion luminosity implied by $\\dot M_\\mathrm{acc}\\simeq5\\times10^{-15}\\Msun\\mathrm{yr^{-1}}$ is $L\\simeq1.8\\times10^{28}\\mathrm{erg\\,s^{-1}}$. Integrating the H$\\alpha$ flux, we find that $F(\\mathrm{H}\\alpha)\\simeq1.0\\times10^{-15}\\mathrm{erg\\,cm^{-2}\\,s^{-1}}$, or, adopting $d=33$\\,pc \\citepalias{tappertetal07-1}, $L(\\mathrm{H}\\alpha)\\simeq1.3\\times10^{26}\\mathrm{erg\\,s^{-1}}$, i.e.~it is clear that the accretion-heated layer of the white dwarf must cool through additional emission mechanisms. In \\citetalias{tappertetal07-1}, we discussed how identifying LTT\\,560 with a nearby faint ROSAT\\,PSPC source would imply an X-ray luminosity of $\\sim6\\times10^{27}\\mathrm{erg\\,s^{-1}}$, which is comparable to the predicted accretion luminosity. A deep X-ray observation would be desirable to confirm the association of LTT\\,560 with the ROSAT X-ray source and to establish more accurately the flux and spectral shape of the X-ray emission.  The nearly identical $\\gamma$ velocities of both the white dwarf's H$\\alpha$ emission and the photospheric metal absorption lines suggest that they originate in the same region. However, because of the uncertainties involved we cannot exclude that the H$\\alpha$ emission is produced somewhat above the white dwarf's photosphere. For a ``worst case scenario\", we calculate the largest possible difference within the 3$\\sigma$ uncertainties as \\begin{displaymath} \\Delta \\gamma = \\gamma_\\mathrm{WD,metal} + 3\\sigma_\\mathrm{metal} - (\\gamma_\\mathrm{WD,H\\alpha}-3\\sigma_\\mathrm{H\\alpha}) = 2.15~\\mathrm{km~s^{-1}}. \\end{displaymath} Since the white dwarf radius is given by $R_\\mathrm{WD} = 0.636~M_\\mathrm{WD}/v_\\mathrm{gr}$ for $M$ and $R$ in solar units and $v_\\mathrm{gr}$ in $\\mathrm{km~s^{-1}}$, $\\Delta \\gamma$ translates into $0.11 R_\\mathrm{WD}$. The H$\\alpha$ emission could thus in principle originate in a region up to $\\sim$1000 km above the photosphere. This still practically excludes a ``strong\" shock as the mechanism behind the emission, because the associated shock height is inversely correlated with the mass-transfer rate \\citep[e.g., Eq.14 in][]{fischer+beuermann01-1}, which for LTT\\,560 is two orders of magnitude below that of even the low-accretion-rate polars \\citep{schwopeetal02-1}. It appears more likely that the emission line is produced by heat deposited by the accretion resulting in a temperature reversal shortly above the white dwarf photosphere. However, a detailed treatment of the physics involved is beyond the scope of this paper.  The narrow metal lines allow us to place a limit on the rotation rate of the white dwarf. Fig.\\,\\ref{f-vsini} shows a close-up of three Fe\\,I lines, along with models for $v\\sin i=0$, 15, and 30\\,km/s, and suggests that the white dwarf in LTT\\,560 is a very slow ($\\la15$\\,km/s) rotator, similar to the majority of single white dwarfs \\citep{koesteretal98-2, karletal05-1, bergeretal05-2}. This is an interesting result, as it suggests that the evolution of angular momentum of white dwarfs in PCEBs is similar to that of field white dwarfs, with current theories requiring magnetic torques to explain the observed low rotation rates \\citep[e.g.][]{suijsetal08-1}.  The data sets used in \\citetalias{tappertetal07-1} and the present one show the occurrence of flares, indicating that there is an active secondary star in the system. Roche tomography shows an asymmetric surface brightness distribution, which we interpret as the presence of star spots (Fig. \\ref{rt_fig}). These high latitude and/or polar spots are very common in single active stars \\citep[e.g., ][]{strassmeieretal03-1}, and have also been observed on donor stars of CVs \\citep{watsonetal06-1,watsonetal07-1}, as well as in the pre-CV V471 Tau \\citep{hussainetal06-1}. The models of \\citet{granzeretal00-1} suggest that star-spots emerge preferentially from the flux tube close to the equator of the star and are dragged by the Coriolis force to the pole of the star. Thus, high latitude and polar spots should be a common feature on rapid rotators, such as close binaries. However, we caution that these models are only valid for stars in the range $0.4~M_\\odot \\leq M \\leq 1.7~M_\\odot$ that contain a radiative core and a convective shell, and in LTT 560 the M5--6V secondary with a mass $M_\\mathrm{MS} = 0.14~M_\\odot$ is well below this range and can be expected to be fully convective.  In addition to the high latitude features, we also observe that the inner face of the star in our model (the one facing the primary) is slightly darker than its backside. The low latitude star-spots populating the region around L1 appear to be another common feature in the tidally distorted late-type secondaries of close binaries \\citep{watsonetal06-1,watsonetal07-1,hussainetal06-1}. It has been suggested by \\citet{holzwarth+schuessler03-1} that tidal interaction may force spots to appear at preferred locations. \\citet{king+cannizzo98-1} find the appearance of such spots at L1 to be the probable mechanism behind the low brightness states and mass-transfer variations in CVs. As \\citet{watsonetal07-1} point out, a similar effect to that of star spots at L1 can be achieved by irradiation from the white dwarf primary, although the low temperature of the white dwarf \\citepalias[$T_\\mathrm{WD} = 7500~\\mathrm{K}$; ][]{tappertetal07-1} indicates that this effect will be small if important at all.  What kind of future awaits LTT 560? Following \\citet{schreiber+gaensicke03-1}\\footnote{As already noted by \\citet{zorotovicetal10-3}, equation (11) in \\citet{schreiber+gaensicke03-1} is incorrect in that the factor $9\\pi$ has to be replaced by $2\\pi$. Since their results have been found to be correct, this is merely a typographical error.}, we calculate the period $P_\\mathrm{sd}$ at which the system becomes semi-detached as \\begin{equation} P_\\mathrm{sd} = 2\\pi \\left(\\frac{R_\\mathrm{MS}^3}{G~M_\\mathrm{MS}~(1+q^{-1}) (R_L/a)^3}\\right)^{\\frac{1}{2}}, \\end{equation} taken from \\citet{ritter86-1}, with the approximation of \\citet{eggleton83-1} \\begin{equation} R_L/a = \\frac{0.49~q^\\frac{2}{3}}{0.6~q^\\frac{2}{3} + \\ln (1+q)^\\frac{1}{3}}. \\end{equation} With the parameters given in Table \\ref{syspars_tab} (taking the average of the He- and the CO-configuration) and the secondary's radius estimated from the RT, we find that $P_\\mathrm{sd} = 1.52^{+0.85}_{-0.43}~\\mathrm{h}$. The large uncertainty here is dominated by our insufficient knowledge of $R_\\mathrm{MS}$. Modelling the light curve of a photometric time-series data set of LTT 560 of high S/N would be desirable to improve the accuracy. Determining the distance to the system via a parallax measurement would similarly yield the absolute luminosities of the stellar components, thus provide independent access to both $R_\\mathrm{MS}$ and $R_\\mathrm{WD}$ and the associated parameters. Nevertheless, a $P_\\mathrm{sd}$ below the period gap is consistent with the M5--6V spectral type of the secondary star \\citep{beuermannetal98-1,smith+dhillon98-1}. Since it can be assumed that the secondary is fully convective, we use Eq.8 from \\citet{schreiber+gaensicke03-1} for angular momentum loss dominated by gravitational radiation to calculate the time it will take LTT 560 to start mass transfer via Roche-lobe overflow to $\\sim$3.5 Gyrs. This is much less than the Hubble time, thus LTT 560 can be regarded as representative of the progenitors of todays CVs. Since the system contains a non-magnetic white dwarf and the mass ratio $q < 0.33$, it is likely that the future CV LTT 560 will belong to the SU UMa subclass of dwarf novae."
    },
    "1107/1107.0470_arXiv.txt": {
        "abstract": "New calculations of the energy levels, radiative transition rates and collisional excitation rates of \\ion{Fe}{ix} have been carried out using the Flexible Atomic Code, paying close attention to experimentally identified levels and extending existing calculations to higher energy levels. For lower levels, R-matrix collisional excitation rates from earlier work have been used. Significant emission is predicted by these calculations in the 5f-3d transitions, which will impact analysis of SDO AIA observations using the 94\\AA\\ filter. ",
        "introduction": "The launch of the Solar Dynamic Observatory (SDO) allows observation of the Sun in unprecedented detail. The Atmospheric Imaging Assembly (AIA, \\citet{lemen2011}) provides multiple simultaneous images of the solar disk every 12 seconds, taken through a variety of narrowband filters centered on individual emission lines of interest. One such filter in centered on 94\\AA, and targets both the \\ion{Fe}{xviii} (93.923\\AA) transition from $2\\mbox{s}^1\\ 2\\mbox{p}^6\\ ^2\\mbox{S}_{1/2}\\rightarrow 2\\mbox{s}^2\\ 2\\mbox{p}^5\\ ^2\\mbox{P}_{3/2}$ at high temperatures and the \\ion{Fe}{x} (94.012\\AA) line from the $3\\mbox{p}^4\\ 4\\mbox{s}^1\\ ^2\\mbox{D}_{5/2}\\rightarrow 3\\mbox{p}^5\\ ^2\\mbox{P}_{3/2}$ transition in cooler plasma. The emission from these two lines occurs in very different temperature ranges and therefore can be distinguished if the plasma temperature distribution is known \\citep{boerner2011}.   \\citet{lepson2002} observed spectra of \\ion{Fe}{vii} to \\ion{Fe}{x} using an electron beam ion trap (EBIT) and a grazing incidence spectrometer with resolution of $\\approx300$ at 100\\AA\\ to observe lines in the 60 to 140\\AA\\ range. They estimated that 70\\% of the emission in this band was unaccounted for by the existing atomic data in the \\textsc{mekal} \\citep{kaastramewe1993, mewe1995} database. Since then, none of the major atomic databases (e.g. \\textsc{chianti} \\citep{dere2009}, AtomDB \\citep{foster2011}, \\textsc{nist} \\citep{ralchenko2011}, and the successor to \\textsc{mekal}, \\textsc{spex} \\citep{kaastra1996}) have updated their atomic data for \\ion{Fe}{ix} to include any lines in this region. Of particular interest are two \\ion{Fe}{ix} lines, observed at 93.59 and 94.07\\AA, which both fall in the bandpass of the 94\\AA\\ AIA filter.  These lines were identified by comparing the EBIT results with calculations using the Hebrew University Lawrence Livermore Atomic Code (\\textsc{hullac}, \\citet{barshalom2001}) as belonging to the $3\\mbox{p}^5 5\\mbox{f}^1 \\rightarrow 3\\mbox{p}^5 3\\mbox{d}^1$ transitions of \\ion{Fe}{ix}. The structure and collisional calculations of the current version of \\textsc{chianti}, version 6.0.1 (and, by extension, AtomDB which uses the \\textsc{chianti} data for this ion) only include configurations up to the n=4 principal quantum shell, and therefore omit these transitions. The NIST \\citep{ralchenko2011} database does include some n=5 transitions but not these lines or levels.  In this work, we have carried out calculations using the Flexible Atomic Code (FAC, \\citet{gu2003}) to extend the energy level and collisional calculations to include higher energy levels up to the n=6 shell, merged the resulting data with the best available collisional and radiative data for lower levels, and used the collisional-radiative code \\textsc{apec} \\citep{smith2001} to model the resulting emission and therefore the effect on the AIA 94\\AA\\ filter flux. ",
        "conclusions": "\\label{sec:discussion} The lack of information on the \\ion{Fe}{ix} emission in the 94\\AA\\ region has been well known since the work of \\citet{lepson2002}. \\citet{testa2011} compared the existing atomic data in \\textsc{chianti} 6.0.1 with observations of Procyon using the Chandra LETG spectrometer, which covers wavelengths up to $\\sim 170$\\AA. They noted in these observations a significant missing flux in the 94\\AA\\ region. Since the launch of the SDO, the lack of a collisional excitation calculation has led to several authors using creative means to account for this difference. When looking at coronal loop emission \\citet{schmelz2011} used a scaling factor obtained by comparing the relative intensities of the known transition lines and the $5\\mbox{f} \\rightarrow 3\\mbox{d}$ in the spectrum observed by \\citet{lepson2002}. They eventually discounted this band from their analysis due to a low count rate, so the effectiveness of this method is unknown. \\citet{aschwanden2011}, when looking at the temperature structure of coronal loops, derived a ``correction factor'', $q_{94}$, for the low temperature ($\\log_{10}(T/K) \\le 6.3$) response function of the 94\\AA\\ filter, $R_{94}$, by fitting their 100 results in the other filters with $q_{94}$ a free parameter. They obtained $q_{94} = 6.7 \\pm 1.7$ at temperatures around $\\log_{10}(T/K)=6.0$. The results from this work imply a much smaller correction factor, with $q_{94}\\approx 2$ being the maximum value at $\\log_{10}(T/K) = 5.8$, and $\\approx 1.25$ at $\\log_{10}(T/K) = 6.0$. There is significant discrepancy here: it would be an interesting exercise to investigate whether the DEM analysis of \\citet{aschwanden2011} can be self consistent using the new emission estimate in this filter band. It is also possible that there are futher lines from different ions in the filter which have not yet been correctly handled.  It is difficult to estimate uncertainties in the results from this work due to the many overlapping sources of error. As well as the two main lines which are the focus of this work, there are  many smaller lines of uncertain wavlength. Given the observed and initially calculated 5f $\\rightarrow$ 3d wavelengths differ by  around 0.5\\%, adding normally distributed random errors to these lines of $\\pm 0.5$\\AA\\ gives a 10\\% standard deviation in the resulting \\ion{Fe}{ix} emission in the 94\\AA\\ filter. This 10\\% fluctuation is then approximately a 5\\% effect on the total emission in the band once \\ion{Fe}{x} is also included. Given that the estimates of uncertainty in distorted wave excitation collisions are usually not better than 20\\%, this is not expected to be a dominating source of error in the calculation.  For those levels which appear in both this work and the collisional calculation of \\citet{storey2002}, comparison of the collision strengths for excitation from the ground state at $T=10^6$K between this data and our FAC results show that they vary by on average around 25\\%. This is a very approximate lower limit on the likely errors on the calculation of emission in the 94\\AA\\ band, since excluding cascades from the n=6 shell, this is the only way to populated the 5f upper levels in this model, and for these levels we are using the distorted wave, not R-matrix, methods.  This data has been incorporated into AtomDB, to be released fully in AtomDB v2.1.0 which is due for release later in 2011. In the meantime, the data can be obtained from the AtomDB website, \\url{www.atomdb.org}."
    },
    "1107/1107.0646_arXiv.txt": {
        "abstract": "The identifications of quasars in the redshift range 2.2$<$z$<$3 are known to be very inefficient as their optical colors are indistinguishable from those of stars. Recent studies have proposed to use optical variability or near-IR colors to improve the identifications of the missing quasars in this redshift range. Here we present a case study by combining both factors. We select a sample of 70 quasar candidates from variables in SDSS Stripe 82, which are non-UV excess sources and have UKIDSS near-IR public data. They are clearly separated into two parts on the Y-K/g-z color-color diagram, and 59 of them meet or lie close to a newly proposed Y-K/g-z selection criterion for z$<$4 quasars. 44 of these 59 sources have been previously identified as quasars in SDSS DR7, and 35 among them are quasars at 2.2$<$z$<$3. We present spectroscopic observations of 14 of 15 remaining quasar candidates using the Bok 2.3m telescope and the MMT 6.5m telescope, and successfully identify all of them as new quasars at z=2.36 to 2.88. We also apply this method to a sample of 643 variable quasar candidates with SDSS-UKIDSS nine-band photometric data selected from 1875 new quasar candidates in SDSS Stripe 82 given by Butler \\& Bloom based on the time-series selections, and find that 188 of them are probably new quasars with photometric redshifts at 2.2$<$z$<$3. Our results indicate that the combination of optical variability and optical/near-IR colors is probably the most efficient way in finding 2.2$<$z$<$3 quasars and very helpful for constructing a complete quasar sample. We discuss its implications to the ongoing and upcoming large optical and near-IR sky surveys. ",
        "introduction": "Since their discovery in 1960s (Schmidt 1963), quasars have become important extra-galactic objects in astrophysics. They not only can be used to probe the physics of supermassive black holes and accretion/jet process, but also are closely related to the studies of galaxy evolution, intergalactic medium, large scale structure and cosmology. More than 120,000 quasars have been discovered from the large optical sky surveys, such as the Two-Degree Fields survey (Boyle et al. 2000) and the Sloan Digital Sky Survey (SDSS, York et al. 2000; Schneider et al. 2010). The quasar candidates in these surveys were mainly selected by optical colors, namely that, due to the strong UV and optical emissions, quasars at $\\rm z<2.2$ and $\\rm z>3.0$ can be distinguished from the stellar objects in the color-color and color-magnitude diagrams based on optical photometry (Smith et al. 2005; Richards et al. 2002; Fan et al. 2000). However, in the redshift range $\\rm 2.2<z<3.0$, the redshifted spectral energy distributions of quasars show similar optical colors to that of normal stars, and quasar selections using the optical color-color diagrams become very inefficient due to the serious contaminations of stars (Fan 1999; Richards et al. 20002; 2006; Scheneider et al. 2007). Because of the crucial importance of z$>$2.2 quasars in studying the Ly$\\alpha$ forest and cosmic baryon acoustic oscillation (BAO) (White 2003; McDonald \\& Eisenstein 2007) and in constructing the accurate luminosity function to study the quasar evolution in the mid-redshift universe (Wolf et al. 2003; Jiang et al. 2006), we have to explore other efficient ways to identify the missing $\\rm 2.2<z<3.0$ quasars.  In the last a few years, two main approaches have been taken to separate quasars and stars rather than using optical color-color diagrams. The first approach is to use optical variability, as this is one of the well known quasar properties (Hook et al. 1994; Cristiani et al. 1996; Giveon et al. 1999). Schmidt et al. (2010) have proposed a method to select quasar candidates by their intrinsic variability. They showed that the quasar structure functions, constructed from the light-curves of known quasars in SDSS Stripe 82 (hearafter S82; see also Sesar et al. 2007), can be modeled by a power-law function with amplitude A and power-law index $\\gamma$. Quasars can be separated from  stars in the $\\rm A-\\gamma$ plane, which enables efficient selection of quasar candidates based on long-term single-band optical photometry (Schmidt et al. 2010). They also pointed out that in the redshift range $\\rm 2.5<z<3.0$, variability can  help to select quasars with a completeness of 90$\\%$. MacLeod et al. (2011) also developed a method to use the damping timescale and asymptotic amplitude of variable sources in S82 to separate quasars from stars with an efficiency higher than 75\\%. Butler \\& Bloom (2011) recently presented a similar time-series study of quasars in S82, and proposed to use two statistics, a quasar-like variability metric and a non-quasar variability metric, to yield the separation of quasar candidates from stars. They claimed that with their method they can achieve nearly a factor of two increase of quasars at $\\rm 2.5<z<3.0$. In addition,  very recent results from the SDSS-III Baryon Oscillation Spectroscopic Survey (BOSS; Eisenstein et al. 2011) also confirmed the high success rate of spectroscopically identifying variability selected quasars, which leads to a significant increase of z$>$2.2 quasar density  in S82 than that based on optical colors only (Palanque-Delabrouille et al. 2011; Ross et al. 2011).   The second approach to separate z$>$2.2 quasars from stars is to utilize their near-IR colors. As the continuum emission from stars usually decreases more rapidly from optical to the near-IR wavelengths than that of quasars, the near-IR colors for stars are different from quasars. This leads to a method to use the K-band excess to identify quasars at z$>$2.2 (eg. Warren, Hewett \\& Foltz 2000; Croom, Warren \\& Glazebrook 2001; Sharp et al. 2002; Hewett et al. 2006; Chiu et al. 2007; Maddox et al. 2008;  Smail et al. 2008; Wu \\& Jia 2010).  Using the photometric data in the $ugriz$ bands of SDSS DR7 (Abazajian et al. 2009) and YJHK bands of UKIRT InfraRed Deep Sky Surveys (UKIDSS\\footnote{ The UKIDSS project is defined in Lawrence et al. (2007). UKIDSS uses the UKIRT Wide Field Camera (WFCAM; Casali et al. 2007) and a photometric system described in Hewett et al. (2006). The pipeline processing and science archive are described in Hambly et al. (2008).}) Large Area Survey (LAS) DR3, Wu \\& Jia (2010) compiled a sample of 8498 SDSS-UKIDSS quasars and a sample of 8996 SDSS-UKIDSS stars. Based on these two samples they compared different optical/near-IR color-color diagrams and proposed an efficient empirical criterion for selecting  z$<$4 quasars in the near-IR Y$-$K and optical g$-$z color-color diagram (i.e. $\\rm Y-K>0.46(g-z)+0.53$, where all magnitudes are Vega magnitudes). With this criterion, they obtained the completeness of 98.6$\\%$ of recovering z$<$4 quasars with the mis-identifying rate of 2.3$\\%$ of classifying stars as quasars. A check with the FIRST (Becker, White \\& Helfand 1995) radio-detected SDSS quasars, which are believed to be free of color selection bias, also proved that with this Y-K/g-z criterion they can achieve the completeness higher than 95$\\%$ for these radio-detected quasars with z$<$3.5, which seems to be difficult in the case of using the SDSS optical color selection criteria alone where a dip around z$\\sim$2.7 in the redshift distribution obviously exists (Richards et al. 2002,2006; Scheneider et al. 2007,2010). Recently, Peth, Ross \\& Schneider (2011) extended the study of Wu \\& Jia (2010) to a much larger sample of 130000 SDSS-UKIDSS selected quasar candidates and re-examined the methods of separating stars and mid-redshift quasars with the near-IR/optical colors. Using the Y-K/g-z selection criterion, Wu et al. (2010a,b) also successfully identified some $\\rm 2.2<z<3.0$ quasars during the commissioning period of the Chinese GuoShouJing Telescope (LAMOST), which provides further supports to the effectiveness of selecting the mid-redshift quasars using the optical/near-IR colors.  Although both approaches we mentioned above can be used to identify quasars at  $\\rm 2.2<z<3.0$, a more ideal approach is to combine the variability and optical/near-IR color to achieve the maximum efficiency. In this paper,  we present a case study by selecting a sample of variable, non-UV excess, SDSS-UKIDSS quasar candidates in S82 (Schmidt et al. 2010), and spectroscopically identifying 14 new quasars at z=2.36 to 2.88. We also apply this method to some new variable quasar candidates in S82 recently suggested by Butler \\& Bloom (2011) and found that 188 SDSS-UKIDSS sources are probably new quasars with 2.2$<$z$<$3. We describe the sample selections and spectroscopic observations in Section 2, present more new  2.2$<$z$<$3 quasar candidates in S82 in Section 3 and discuss the results in Section 4. ",
        "conclusions": "We have presented a case study to demonstrate that we can effectively select  $\\rm 2.2<z<3.0$ quasars by combining the optical variability and optical/near-IR colors. Our successful spectroscopic identifications of 14 new quasars at z=2.36 to 2.88 with the Bok 2.3m telescope and the MMT 6.5m telescope provide further support to this combination approach, which can be used to select quasars with probably the highest efficiency (here we define the efficiency as the percentage of quasars identified from the spectroscopic targets, similar to the definition in SDSS-III (Ross et al. 2011)). We also compiled a catalog of 188 quasar candidates with photometric redshifts at $\\rm 2.2<z<3.0$ from variable SDSS-UKIDSS sources in S82, and expect that the ongoing SDSS III spectroscopy will confirm their quasar nature and redshifts soon.  We noticed that although combining the optical/near-IR colors and time-series information can help increase the efficiency in identifying quasars, it may decrease the completeness of quasars if both selection criteria on colors and variability are required. This can also be seen from Fig. 1 and Fig. 4 as some quasars selected by variability do not meet our color selection criterion. One possible way to avoid this is to decrease the threshold for each criterion. For example, relaxing our color criterion to $Y-K>0.46(g-z)+0.50$ would include all variability-selected quasars in Fig. 1 and 98.8\\% of variability-selected quasar candidates in Fig. 4, without increasing much contamination from stars as most of them are less variable and still far away from our color selection criterion. However, how to best combine both the optical/near-IR colors and time-series information to select quasars with both higher efficiency and higher completeness obviously needs more investigations in the future based on complete samples of quasars and stars in certain sky areas.  For using this combination approach, we need both the optical variability measurements and optical/near-IR photometry in a large sky area. Especially for finding z$>2.2$ quasars, deeper imaging and multi-epoch photometry are neccesary. However, so far both variability and optical/near-IR photometric observations have been realized only for a small part of the sky, such as in S82. Because the typical variability timescales of quasars are usually in years in the optical band, we need to measure the variability of sources for many epochs in at least several years in order to get better statistics to determine their variability features. That is why so far the variability studies related to quasars have been done only on some smaller sky areas, which significantly limits the efforts in discovering quasars with variability. However, even if we may not have both time-series and color information for quasar candidate selection in most sky areas,  utilizing both information as much as possible often allows us to get the most quasars. Fortunately, there are also several ongoing and upcoming large projects with both photometric and variability information, especially the Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS; Kaiser et al. 2002) and the Large Synoptic Survey Telescope (LSST; Ivezic et al. 2008). The multi-epoch photometry in multi-bands covering a large part of the sky by these facilities will hopefully provide better opportunities to use variability to construct a much larger sample of quasars than that currently available.  On the other hand, several ongoing and upcoming optical and near-IR photometric sky surveys will also provide crucial helps to us to extend the SDSS-UKIDSS optical/near-IR color selection of quasars to larger and deeper fields. In addition to SDSS III (Eisenstein et al. 2011), which has taken 2500 deg$^2$ further imaging in the south galactic cap, the SkyMapper (Keller et al. 2007) and Dark Energy Survey (DES; The Dark Energy Survey Collaboration 2005) will also present the multi-band optical photometry in 20000/5000 deg$^2$ of the southern sky, with the magnitude limit of 22/24 mag in $i$-band, respectively. The Visible and Infrared Survey Telescope for Astronomy (VISTA; Arnaboldi et al. 2007) will carry out its VISTA Hemisphere Survey (VHS) in the near-IR YJHK bands for 20000 deg$^2$ of the southern sky with a magnitude limit at K=20.0, which is about five magnitude and two magnitude deeper than the Two Micron ALL Sky Survey (2MASS; Skrutskie et al. 2006) and UKIDSS/LAS limits (Lawrence et al. 2007), respectively. Therefore, the optical and near-IR photometric data obtained with these ongoing and upcoming surveys will provide us a large database for quasar selections. By combining the optical variability and the optical/near-IR colors, we expect that a much larger and more complete quasar sample can be efficiently constructed in the near future.  Although by combining the variability and optical/near-IR colors we can efficiently select quasar candidates and reliably estimate their photometric redshifts, the spectroscopic identifications are still crucial to determine their quasar nature and redshifts. The ongoing BOSS project in SDSS III has identified 29000 quasars with $z>2.2$ and expects to obtain the spectra of 150000 quasars at  $2.2<z<4$ (Eisenstein et al. 2011; Ross et al. 2011). We believe that many $\\rm 2.2<z<3.0$ quasars, including the candidates we listed in this paper,  should be spectroscopically identified by BOSS. In addition, the Chinese GuoShouJing telescope (LAMOST; Su et al. 1998), a spectroscopic telescope with 4000 fibers currently in the commissioning phase, is also aiming at discovering 0.3 million quasars with magnitudes bright than $i=20.5$ (Wu et al. 2010b). By combining the variability and optical/near-IR colors, large input catalogs of reliable quasar candidates will be provided to these quasar surveys for future spectroscopic observations. Therefore, we expect that a much larger and more complete quasar sample covering a wider range of redshift  will be constructed in the near future, which will play an important role in studying extra-galactic astrophysics, including  the physics of accretion around supermassive black holes, galaxy evolution, intergalactic medium, large scale structure and cosmology."
    },
    "1107/1107.5313_arXiv.txt": {
        "abstract": "{NGC7538 is a complex massive star-forming region. The region is composed of several radio continuum sources, one of which is IRS\\,1, a high-mass protostar, from which a 0.3~pc molecular bipolar outflow was detected. Several maser species have been detected around IRS\\,1. The \\meth ~masers have been suggested to trace a Keplerian-disk, while the \\hdueo ~masers are almost aligned to the outflow. More recent results suggested that the region hosts a torus and potentially a disk, but with a different inclination than the Keplerian-disk that is supposed to be traced by the \\meth ~masers.} {Tracing the magnetic field close to protostars is fundamental for determining the orientation of the disk/torus. Recent studies showed that during the protostellar phase of high-mass star formation the magnetic field is oriented along the outflows and around or on the surfaces of the disk/torus. The observations of polarized maser emissions at milliarcsecond resolution can make a crucial contribution to understanding the orientation of the magnetic field and, consequently, the orientation of the disk/torus in NGC7538-IRS\\,1.} {The NRAO Very Long Baseline Array was used to measure the linear polarization and the Zeeman-splitting of the 22\\,GHz \\hdueo ~masers toward NGC7538-IRS\\,1. The European VLBI Network and the MERLIN telescopes were used to measure the linear polarization and the Zeeman-splitting of the 6.7\\,GHz \\meth ~masers toward the same region.} {We detected 17 \\hdueo ~masers and 49 \\meth ~masers at high angular resolution. We detected linear polarization emission toward two \\hdueo ~masers and toward twenty \\meth ~masers. The \\meth ~masers, most of which only show a core structure, seem to trace rotating and potentially infalling gas in the inner part of a torus. Significant Zeeman-splitting was measured in three \\meth ~masers. No significant (3$\\sigma$) magnetic field strength was measured using the \\hdueo ~masers. We also propose a new description of the structure of the NGC7538-IRS\\,1 maser region.} {} ",
        "introduction": "\\indent NGC7538 is a complex massive star-forming region located in the Perseus arm of our Galaxy at a distance of 2.65 kpc (Moscadelli et al. \\cite{mos09}). The region is composed of several clusters of infrared sources (Wynn-Williams et al. \\cite{wyn74}) and radio continuum sources (Campbell \\cite{cam84}). The brightest source is NGC7538-IRS\\,1, whose central star has been suggested to be an O6 star of about 30~\\solmass ~with systemic local standard of rest velocity $V_{\\rm{lsr}}=-58$~\\kms ~(Campbell \\& Thompson \\cite{cam84b}; Sandell et al. \\cite{san09}; Puga et al. \\cite{pug10}). Several high-velocity molecular bipolar outflows were detected in NGC7538 and one of these is elongated 0.3~pc from IRS\\,1 (position angle PA=140\\d) with a velocity of 250~\\kms ~and a mass of 82.8~\\solmass ~(Kameya et al. \\cite{kam89}; Gaume et al. \\cite{gau95}; Davis et al. \\cite{dav98}; Qiu et al. \\cite{qiu11}). VLA continuum observations by Campbell (\\cite{cam84}) indicate that the PA of the outflow decreases away from IRS\\,1; i.e., 180\\d ~at 0$\\farcs$3 (0.004~pc) and 165\\d ~at $2''$ (0.03~pc). Kameya et al. (\\cite{kam89}) gave three possible interpretations of this rotation: disk precession, interaction of the flow with dense gas, and coupling of the gas with a large-scale magnetic field around IRS\\,1. Because Sandell et al. (\\cite{san09}) found that the collimated free-free jet (opening angle $\\lesssim 30$\\d) is approximately aligned with the outflow and that there is a strong accretion flow toward IRS\\,1 (accretion rate $\\sim2\\times10^{-4}$~\\solmass$\\rm{yr^{-1}}$), IRS\\,1 must be surrounded by an accretion disk. The morphology of the free-free emission, which is optically thick up to 100~GHz (Franco-Hern\\'{a}ndez  \\& Rodr\\'{i}guez \\cite{fra04}), suggests that the disk should be almost edge-on and oriented east-west (Scoville et al \\cite{sco86}; Kameya et al \\cite{kam89}; Sandell et al. \\cite{san09}). A possible detection of this edge-on disk was made by Minier et al. (\\cite{min98}), who observed a linear distribution of the brightest \\meth ~masers with an inclination angle of about 112\\d. Pestalozzi et al. (\\cite{pes04}) estimated that this disk is a Keplerian disk with an outer radius of $\\sim$750~AU and an inner radius of $\\sim$290~AU by modeling the \\meth ~maser emissions at 6.7 and 12.2-GHz.\\\\ \\indent Recent results disagree with the edge-on disk traced by the \\meth ~masers. The observations of the mid-infrared emission suggest that the radio continuum emission does not come from a free-free jet but that it traces the ionized gas wind from the disk surface that in the new scenario would be perpendicular to the CO-bipolar outflow with a disk inclination angle $i=32$\\d~ (De~Buizer \\& Minier \\cite{bui05}). Klaassen et al. (\\cite{kla09}) observed two warm gas tracers (SO$_{2}$ and OCS) toward NGC7538-IRS\\,1 with the Submillimeter Array (SMA). Although the region was unresolved, they found a velocity gradient consistent with the \\meth ~maser velocities and perpendicular to the large-scale molecular bipolar outflow. This rotating gas might indicate that there is a torus (with an angular size of about $2$ arcsec; i.e., $\\sim$5300~AU at 2.65~kpc) surrounding the smaller accretion disk proposed by De Buizer \\& Minier (\\cite{bui05}).\\\\ \\indent In addition, Krauss et al. (\\cite{kra06}) proposed two possible scenarios for NGC7538-IRS\\,1 by considering both the linear distribution of the 6.7-GHz \\meth ~masers detected previously (e.g., Pestalozzi et al. \\cite{pes04}, De Buizer \\& Minier \\cite{bui05}) and the asymmetry in the near infrared (NIR) images observed by them. In their most likely scenario, called ``Scenario B'', the \\meth ~masers trace the edge-on disk as suggested by Pestalozzi et al. (\\cite{pes04}) and the asymmetry might be caused by the precession of the jet. In the ``Scenario A'' the \\meth ~masers do not trace a disk but an outflow cavity as proposed by De Buizer \\& Minier (\\cite{bui05}), where the detected asymmetry might simply reflect the innermost walls of this cavity.\\\\ \\begin{figure*}[th!] \\centering \\includegraphics[width = 13 cm]{NGC7538_final5.eps} \\caption{Positions of water, methanol, and hydroxyl maser features superimposed on the 2\\,cm continuum contour map of NGC7538-IRS\\,1 observed with the VLA in 1986 (Franco-Hern\\'{a}ndez  \\& Rodr\\'{i}guez \\cite{fra04}). Contours are 1, 2, 4, 8, 16 $\\times$1.3 \\mjyb. The (0,0) position is $\\alpha_{2000}=23^{\\rm{h}}13^{\\rm{m}}45^{\\rm{s}}\\!.382$ and $\\delta_{2000}=+61^{\\circ}28'10''\\!\\!.441$. Blue circles and red triangles indicate the positions of the \\hdueo ~and \\meth ~maser features, respectively, with their linear polarization vectors (60 mas correspond to a linear polarization fraction of 1\\%). Green boxes indicate the positions of the OH masers detected by Hutawarakorn \\& Cohen (\\cite{hut03}) with their linear polarization vectors (thick lines, 6~mas correspond to 1\\%). The two arrows indicate the direction of the bipolar outflow ($\\sim$140\\d).} \\label{pos} \\end{figure*} \\indent Besides the \\meth ~masers, other maser species were detected around NGC7538-IRS\\,1: OH, \\hdueo, NH$_{3}$, and H$_{2}$CO (e.g., Hutawarakorn \\& Cohen \\cite{hut03}; Galv\\'{a}n-Madrid et al. \\cite{gal10}; Gaume et al. \\cite{gau91}; Hoffman et al. \\cite{hof03}). The OH masers are located southward and show no obvious disk structure or relation to the outflow direction (Hutawarakorn \\& Cohen \\cite{hut03}). H$_{2}$CO and \\hdueo ~masers are located near the center of the continuum emission, and the \\hdueo ~masers are also almost aligned with the outflow (Galv\\'{a}n-Madrid et al. \\cite{gal10}). The 6.7 and 12.2-GHz \\meth ~masers show a cone shape that opens to the north-west (Minier et al. \\cite{min00}). \\\\ \\indent So far, the magnetic field structure in NGC7538-IRS\\,1 has been studied using submillimeter imaging polarimetry (Momose et al. \\cite{mom01}) and OH maser emission (Hutawarakorn \\& Cohen \\cite{hut03}). Although the polarization vectors are locally disturbed, at an angular resolution of 14~arcsec the magnetic field directions agree with the direction of the outflow (Momose et al. \\cite{mom01}), while at milliarcsecond (mas) resolution the OH maser observations indicate a magnetic field oriented orthogonal to the outflow (Hutawarakorn \\& Cohen \\cite{hut03}). Because \\hdueo ~and \\meth ~masers were detected close to the center of the continuum emission, i.e. to the protostar, it is worthwhile investigating their linear and circular polarization emissions. As shown by Surcis et al. (\\cite{sur09}, \\cite{sur11}) for the massive star-forming region W75N, the polarization observations of the two maser species can depict a reasonable scenario for the magnetic field close to massive protostars. Moreover, the direction and strength of the magnetic field might help to decide the debate about the orientation of the disk.\\\\ \\indent Here we present Very Long Baseline Array (VLBA) observations of \\hdueo ~masers, Multi-Element Radio Linked Interferometer network (MERLIN) and European VLBI Network (EVN) observations of \\meth~ masers in full polarization toward NGC7538-IRS\\,1. In Sect.~\\ref{res} we show our results obtained by studying the linear and circular polarization of \\hdueo ~and \\meth ~maser emissions in a similar way as Vlemmings et al. (\\cite{vle06a}, \\cite{vle10}) and Surcis et al. (\\cite{sur09}, \\cite{sur11}) for Cepheus~A and W75N, respectively. In Sect.~\\ref{dis} we discuss our results and attempt to disentangle the complex morphology in NGC7538-IRS\\,1. \\begin {table*}[th!] \\caption []{All 22 GHz \\hdueo ~maser features detected in NGC7538-IRS\\,1.} \\begin{center} \\scriptsize \\begin{tabular}{ l c c c c c c c c c c c c c} \\hline \\hline \\,\\,\\,\\,\\,(1)&(2)& (3)     & (4)       & (5)           & (6)            & (7)                 & (8)         & (9)        & (10)                    & (11)                    & (12)            & (13)   \\\\ Maser    & Group & RA      & Dec       & Peak flux     & $V_{\\rm{lsr}}$ & $\\Delta v\\rm{_{L}}$ &$P_{\\rm{l}}$ &  $\\chi$    & $\\Delta V_{\\rm{i}}^{a}$ & $T_{\\rm{b}}\\Delta\\Omega^{a}$& $P_{\\rm{V}}^{b}$& $\\theta^{c}$\\\\ &       &offset   & offset    & Density(I)    &                &                     &             &\t         &                         &                         &                 &              \\\\ &       &(mas)    & (mas)     & (\\jyb)        &  (\\kms)        &      (\\kms)         & (\\%)        &   (\\d)     & (\\kms)                  & (log K sr)              & (\\%)    & (\\d)       \\\\ \\hline W01\t & N     &-185.2131& 141.908   & $17.11\\pm0.07$& -57.8          & \t0.41\t       &$-$          & $-$        & $-$                     & $-$                     &\t$-$\t   & $-$ \\\\ W02\t & N     &-176.4676& 139.422   & $1.00\\pm0.02$ & -58.3          & \t0.32\t       &$-$          & $-$        & $-$ \t\t    & $-$\t              &\t$-$\t   & $-$ \\\\ W03\t & N     &-170.9744& 139.088   & $0.48\\pm0.02$ & -58.3          & \t0.27\t       &$-$          & $-$        & $-$     \t\t    & $-$\t     \t      &\t$-$\t   & $-$         \\\\ W04\t & S     &-57.6929 & -1071.104 & $0.90\\pm0.01$ & -72.2          & \t0.59\t       &$-$          & $-$        & $-$                     & $-$                     &\t$-$\t   & $-$ \\\\ W05\t & S     &-57.4470 & -1056.554 & $0.48\\pm0.01$ & -57.3          &\t0.47\t       &$-$          & $-$        & $-$                     & $-$                     &\t$-$\t   & $-$ \\\\ W06\t & S     &-57.3376 & -1056.630 & $36.43\\pm0.02$& -70.3          &\t1.15\t       &$1.2\\pm0.1$  & $-13.0\\pm1.3$& $-$                   & $-$                     &\t$-$\t   & $-$ \\\\ W07\t & S     &-57.1463 & -1057.065 & $16.99\\pm0.01$& -73.3          & \t0.77\t       &$1.2\\pm0.1$  & $-10.8\\pm0.5$& $3.4^{+0.1}_{-0.3}$   & $9.1^{+0.3}_{-0.1}$     &\t$0.37\\pm0.13$& $81^{+9}_{-13}$\\\\ W08\t & S     &-56.7637 & -1056.586 & $0.15\\pm0.01$ & -67.3          & \t1.46\t       &$-$          & $-$        & $-$                     & $-$     \t\t      &\t$-$\t   & $-$ \\\\ W09\t & S     &-56.4631 & -1070.080 & $0.48\\pm0.01$ & -71.6          & \t0.64\t       &$-$  \t     & $-$\t  & $-$\t\t            & $-$     \t\t      &\t$-$\t   & $-$ \\\\ W10\t & N     &0        & 0         & $7.34\\pm0.03$ & -60.6          & \t0.56\t       &$-$          & $-$        & $-$\t\t            & $-$ \t\t      &\t$-$\t   & $-$ \\\\ W11\t & N     &85.5692 & -96.002    & $0.36\\pm0.02$ & -59.7          & \t0.44\t       &$-$  \t     & $-$\t  & $-$\t\t            & $-$ \t\t      &\t$-$\t   & $-$ \\\\ W12\t & N     &110.0019 & -13.538   & $0.20\\pm0.03$ & -60.2          & \t0.83\t       &$-$  \t     & $-$ \t  & $-$                     & $-$  \t\t      &\t$-$\t   & $-$ \\\\ W13\t & N     &111.6690 & -99.499   & $9.31\\pm0.03$ & -60.1 \t        & \t0.47\t       &$-$  \t     & $-$        & $-$\t\t            & $-$  \t\t      &\t$-$\t   & $-$ \\\\ W14\t & N     &112.8442 & -156.492  & $0.25\\pm0.03$ & -60.1          &\t0.42\t       &$-$  \t     & $-$\t  & $-$   \t\t    & $-$     \t\t      &\t$-$\t   & $-$ \\\\ W15\t & N     &151.9803 & -90.046   & $1.04\\pm0.03$ & -59.9          &\t0.40\t       &$-$ \t     & $-$\t  & $-$ \t\t    & $-$   \t\t      &\t$-$\t   & $-$ \\\\ W16\t & N     &242.0042 & 148.947   & $0.13\\pm0.01$ & -43.3\t\t& \t0.62\t       &$-$  \t     & $-$\t  & $-$\t\t            & $-$\t\t      &\t$-$\t   & $-$ \\\\ W17\t & N     &242.7694 & 149.715   & $0.14\\pm0.01$ & -43.8       \t& \t0.59\t       &$-$  \t     & $-$\t  & $-$\t\t            & $-$    \t\t      &\t$-$\t   & $-$ \\\\ \\hline \\end{tabular} \\end{center} \\scriptsize{\\textbf{Notes.} $^{(a)}$ The best-fitting results obtained by using a model based on the radiative transfer theory of \\hdueo ~masers for $\\Gamma+\\Gamma_{\\nu}=1$s$^{-1}$ (Surcis et al. \\cite{sur11}). The errors were determined by analyzing the full probability distribution function. $^{(b)}$ The percentage of circular polarization is given by $P_{\\rm{V}}=(V_{\\rm{max}}-V_{\\rm{min}})/I_{\\rm{max}}$. $^{(c)}$The angle between the magnetic field and the maser propagation direction is determined by using the observed $P_{\\rm{l}}$ and the fitted emerging brightness temperature. The errors were determined by analyzing the full probability distribution function.} \\label{poltw} \\end{table*} ",
        "conclusions": "The massive star-forming region NGC7538-IRS\\,1 has been observed at 22-GHz in full polarization spectral mode with the VLBA and at 6.7-GHz with the EVN and MERLIN to detect linear and circular polarization emission from \\hdueo ~and \\meth ~masers, respectively. We detected 17 \\hdueo ~masers and 49 \\meth ~masers at high angular resolution. We have measured Zeeman-splitting for three \\meth ~masers ranging from -2.7~\\ms ~to +2.7~\\ms. No significant magnetic field strength has been measured from the \\hdueo ~masers. Furthermore, we have also shown that the masers of NGC7538-IRS\\,1 are all consistent with a torus-outflow scenario. Here the \\meth ~masers are tracing the interface between the infall and the large-scale torus, and the \\hdueo ~masers are related to the blue-shifted part of the outflow. The \\hdueo ~masers of the southern group are thought to be associated with another source. The magnetic field is situated on the two surfaces of the torus with a counterclockwise direction on the top surface.\\\\  \\noindent\\small{\\textit{Acknowledgments.} We wish to thank an anonymous referee for making useful suggestions that have improved the paper. GS, WHTV, and RMT acknowledge support by the Deutsche Forschungsgemeinschaft (DFG) through the Emmy Noether Reseach grant VL 61/3-1. G.S. and WHTV thank Pamela D. Klaassen for the very useful discussion. G.S. thanks Ramiro Franco-Hern\\'{a}ndez for kindly providing the VLA continuum image.}"
    },
    "1107/1107.4253_arXiv.txt": {
        "abstract": "Data gaps are ubiquitous in spectral irradiance data, and yet, little effort has been put into finding robust methods for filling them. We introduce a data-adaptive and nonparametric method that allows us to fill data gaps in multi-wavelength or in multichannel records. This method, which is based on the iterative singular value decomposition, uses the coherency between simultaneous measurements at different wavelengths (or between different proxies) to fill the missing data in a self-consistent way. The interpolation is improved by handling different time scales separately.  Two major assets of this method are its simplicity, with few tuneable parameters, and its robustness. Two examples of missing data are given: one from solar EUV observations, and one from solar proxy data. The method is also appropriate for building a composite out of partly overlapping records. ",
        "introduction": "Solar and stellar irradiance records are often plagued by data gaps. The proper interpolation of these missing data is a longstanding and notoriously delicate problem that requires a good understanding of the data \\citep{wiener64,little02}. Considerable attention has been given to this problem in fields such as climate science \\citep{dobesch07} but much less so in solar physics and in astrophysics. Often, the limited attention that is paid to data gaps contrasts with the sophistication of the analysis that is subsequently performed on these data.  While short gaps can easily be filled by linear or by nonlinear interpolation, data gaps whose duration exceeds the characteristic time scales are much more difficult to handle. A notable exception is when multichannel synoptic observations of the same process are available, with gaps in some or in all of them. Spectral irradiance observations, which we shall concentrate on, precisely belong to that category. Our examples will be taken from the Sun, but the results can be easily extended to other types of multichannel observations. Our method applies to any set of observations that are recorded simultaneously (i.e. the time stamps are the same for all records), are correlated with each other, and whose time intervals fully or partly overlap. Our main assumption is their linear correlation, in the sense that each record can be approximated by a linear combination of the other ones.  A strong linear correlation is typically observed between spectral irradiance observations made at different wavelengths or between simultaneous measurements of different proxies. These synoptic records are frequently used to assess subtle changes in the variability of the Sun; they are often remarkably coherent in time $t$ and in wavelength $\\lambda$. As a consequence, their variability can be explained in terms of a few contributions only. This property  is well known for the Extreme UltraViolet (EUV)  \\citep{lean82,amblard08} but also for the visible range \\citep{rabbette01}, when measured from space. The same coherency is observed among different proxies for solar activity \\citep{pap92b,schmahl94,lean00,kane02c,floyd05,ddw09}. This property is rooted in the  structuring effect of the solar magnetic field. The coherency partly breaks down during the impulsive phase of solar flares because the spectrum then considerably depends on the local conditions of the solar atmosphere. Here, however, as in many applications, we consider daily or hourly averages, so that the effect of short transients can be discarded.  This coherency in both time and wavelength is the key to the reconstruction technique we shall introduce below. By interpolating along two dimensions (in time and along different records), we not only improve the quality of the reconstruction, but we also can fill arbitrarily large data gaps without having to rely on the tedious bookkeeping that is required by most interpolation schemes.  The nonparametric and data-adaptive method we advocate is based on the SVD or singular value decomposition \\citep{golub00}, which is to linear algebra what the Fourier transform is to spectral analysis. The SVD allows the extraction of the coherent part of the solar spectral irradiance, which is then used to fill the data gaps iteratively. The method is described in Section 2, and two applications are detailed. The first one (Sec. 3) deals with solar spectral irradiance data in the EUV. In the second application (Sec. 4) we consider a set of solar proxies with numerous gaps. ",
        "conclusions": "This study shows that SVD-interpolation is a powerful technique for filling arbitrarily large gaps in multi-wavelength, multichannel or in synoptic records. We focused here on solar spectral irradiance observations, which are frequently plagued by missing data.  These gaps may be distributed at random in time or in wavelength. The main tuneable parameter is the number of SVD modes that is needed to reconstruct the data; this value may be estimated either by visualisation or by cross-validation. The method works best when each record can be approximated by a linear combination of the others. Since it relies on linear combinations only, it may be desirable to apply  a nonlinear static transform beforehand to increase the linear correlation between the records.  For the method to work, the observations must be sampled simultaneously but not necessarily evenly. Non-simultaneous observations can be handled by resampling all variables to a common grid, for example by Fourier decomposition \\citep[e.g.][]{hocke09}, and then filling the gaps by SVD. By alternating between the two, both the gaps and the interpolated values can be progressively refined.  This method has several applications in addition to mere interpolation. The first one is the cross-calibration of measurements of the same quantity by different instruments. The Mg II index, for example, is at present measured by different instruments that give different amplitudes. These  data sets are incomplete and only partly overlap, which considerably impairs their inter-comparison. The iterative SVD method is ideally suited for filling these gaps because the records are by definition strongly correlated.  A second potential application is the stitching together of total solar irradiance (TSI) observations. Merging TSI records from several instruments is a delicate and controversial task \\citep{froehlich02} because instruments disagree on the absolute value of the TSI and often do not operate simultaneously. The iterative SVD provides a means for estimating the different offsets in a self-consistent way because it allows us to extrapolate each TSI record by assuming that its statistical properties with respect to the other records do not change in time. This property is particularly useful for checking composites that are built from different records, such as the TSI, the H I Lyman $\\alpha$ intensity, the Mg II index and the sunspot index \\citep{clette07}. This will be detailed in a forthcoming publication.            \\subsubsection*{Acknowledgements} {\\small I thank the following institutes for providing the data: the Laboratory for Atmospheric and Space Physics (University of Colorado) for the Mg II and H I Lyman $\\alpha$ composites, the National Solar Observatory at Sacramento Peak (data produced cooperatively by NSF/NOAO, NASA/GSFC and NOAA/SEC) for the Ca K index, the Mount Wilson Observatory (operated by UCLA, with funding from NASA, ONR  and NSF, under agreement with the Mt. Wilson Institute) for the MPSI index and the Space Sciences Center (University of Southern California) for the SEM data.  This study received funding from the European Community's Seventh Framework Programme (FP7/2007-2013) under the grant agreement nr. 218816 (SOTERIA project, http://www.soteria-space.eu).}"
    },
    "1107/1107.1689_arXiv.txt": {
        "abstract": "Are giant flares in magnetars viable sources of gravitational radiation? Few theoretical studies have been concerned with this problem, with the small number using either highly idealized models or assuming a magnetic field orders of magnitude beyond what is supported by observations. We perform nonlinear general-relativistic magnetohydrodynamics simulations of large-scale hydromagnetic instabilities in magnetar models. We utilise these models to find gravitational wave emissions over a wide range of energies, from $10^{40}$ to $10^{47} \\, \\mathrm{erg}$. This allows us to derive a systematic relationship between the surface field strength and the gravitational wave strain, which we find to be highly nonlinear. In particular, for typical magnetar fields of a few times $10^{15} \\, \\mathrm{G}$, we conclude that a direct observation of $f$-modes excited by global magnetic field reconfigurations is unlikely with present or near-future gravitational wave observatories, though we also discuss the possibility that modes in a low-frequency band up to $100 \\, \\mathrm{Hz}$ could be sufficiently excited to be relevant for observation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.3159_arXiv.txt": {
        "abstract": "We present the results of our ultra-deep Keck/DEIMOS spectroscopy of $z$-dropout galaxies in the SDF and GOODS-N. For $3$ out of $11$ objects, we detect an emission line at $\\sim 1\\mu$m with a signal-to-noise ratio of $\\sim 10$. The lines show asymmetric profiles with high weighted skewness values, consistent with being Ly$\\alpha$, yielding redshifts of $z=7.213$, $6.965$, and $6.844$. Specifically, we confirm the $z=7.213$ object in two independent DEIMOS runs with different spectroscopic configurations. The $z=6.965$ object is a known Ly$\\alpha$ emitter, IOK-1, for which our improved spectrum at a higher resolution yields a robust skewness measurement. The three $z$-dropouts have Ly$\\alpha$ fluxes of $3 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and rest-frame equivalent widths EW$_0^{{\\rm Ly}\\alpha} = 33-43${\\AA}. Based on the largest spectroscopic sample of $43$ $z$-dropouts that is the combination of our and previous data, we find that the fraction of Ly$\\alpha$-emitting galaxies (EW$_0^{{\\rm Ly}\\alpha} > 25${\\AA}) is low at $z \\sim 7$; $17 \\pm 10${\\%} and $24 \\pm 12${\\%} for bright ($M_{\\rm UV} \\simeq -21$) and faint ($M_{\\rm UV} \\simeq -19.5$) galaxies, respectively. The fractions of Ly$\\alpha$-emitting galaxies drop from $z \\sim 6$ to $7$ and the amplitude of the drop is larger for faint galaxies than for bright galaxies. These two pieces of evidence would indicate that the neutral hydrogen fraction of the IGM increases from $z \\sim 6$ to $7$, and that the reionization proceeds from high- to low-density environments, as suggested by an inside-out reionization model. ",
        "introduction": "\\label{sec:introduction}  Over the last two years, we have witnessed an explosion of activities aimed at searching for very high redshift ($z > 7$) galaxies. This has been made possible by the installation of the Wide-Field Camera 3 (WFC3) on-board the Hubble Space Telescope \\citep[e.g.,][]{oesch2010,oesch2011,bouwens2010,bouwens2011,bouwens2011b,mclure2009,mclure2011,wilkins2010,wilkins2011,bunker2010,yan2010,yan2010b,lorenzoni2010}, the Hawk-I instrument on the Very Large Telescope \\citep[VLT;][]{castellano2010,castellano2010b}, and improvement in the red-end sensitivity of the Suprime-Cam on the Subaru telescope \\citep{ouchi2009b}. These studies have identified over a hundred candidates at $z > 6.5$ through the Lyman break drop-out technique \\citep[e.g.,][]{steidel1996a,giavalisco2002}, showing a decrease in the number density of bright high-redshift galaxies with redshift \\citep[e.g.,][]{ouchi2009b,mclure2009,castellano2010,bouwens2010,wilkins2010}, bluer UV continua \\citep[e.g.,][]{bouwens2010b}\\footnote{See also \\cite{dunlop2011}, who question this result.}, and relatively smaller stellar masses compared to lower redshift galaxies selected by similar techniques \\citep[e.g.,][]{labbe2010,schaerer2010,finkelstein2009f,ono2010}. Study of the intrinsic properties of high-redshift galaxies at epochs close to the dark ages is essential for understanding one of the most outstanding questions in modern astronomy - when did the Universe become reionized and what sources were responsible for it? This can be accomplished by studying the state of the inter-galactic medium (IGM) through estimates of the ionizing photon budget and Lyman $\\alpha$ escape fraction.    Several independent studies of the reionization process in recent years have yielded different results. By measuring the polarization of the cosmic background radiation, \\cite{dunkley2009} estimated the optical depth to reionization and concluded that, if it was a sudden event, reionization occurred at $z = 11.0 \\pm 1.4$  \\citep[see also,][]{komatsu2010,larson2011}. However, investigating the spectra of SDSS quasars, \\cite {fan2006} studied the evolution of the Gunn-Peterson optical depth and demonstrated that the IGM reionization may have ended as late as $z \\sim 6$. The result from this study is questioned by \\cite{goto2011c}, who have shown that for quasars at $z > 6$, a statistically large number is needed in $\\Delta z = 0.1$ bins to constrain cosmic variance and trace the evolution of the optical depth.    Another useful tool for studying reionization is the Lyman $\\alpha$ luminosity function of high-redshift galaxies selected by narrow-band imaging \\citep[e.g.,][]{hu1998,rhoads2000}; these galaxies are called Lyman Alpha Emitters (LAEs). Since neutral hydrogen in the IGM resonantly scatters Ly$\\alpha$ photons, the transmission of Ly$\\alpha$ is sensitive to the ionization state of the IGM. Therefore, we expect a decrease in the number density of LAEs close to the reionization epoch \\citep[e.g.,][]{haiman1999,malhotra2004,santos2004,mesinger2004,stern2005,haiman2005,furlanetto2006,mcquinn2007,dijkstra2007,kobayashi2007,mesinger2008b,iliev2008,dayal2008,dayal2009,dayal2011}. \\cite{ouchi2010} found a decrease in the Ly$\\alpha$ luminosity function of LAEs, corresponding to an upper limit of the IGM neutral fraction $x_{\\rm HI} \\lesssim 0.2 \\pm 0.2$ at $z=6.6$, which indicates that the major reionization process took place at higher redshift. This result has been further supported by \\cite{kashikawa2011}, who studied the evolution of the LAE luminosity function in the Subaru Deep Field \\citep[SDF;][]{kashikawa2004} and found an increase in the neutral fraction of the IGM from $z=5.7$ to $6.5$. \\cite{nakamura2011} have also found a large deficit in the number density of $z=6.5$ LAEs in the SSA22 field, and attributed it to significant field-to-field variance of the neutral fraction of the IGM. The effect of cosmic variance was quantified by \\cite{ouchi2010}, who claimed a factor of $2-10$ range in the number density of LAEs in an extensive survey covering an area of $1$ deg$^2$.    \\begin{deluxetable*}{cccccccc} \\tablecolumns{8} \\tablewidth{0pt} \\tablecaption{Summary of observations with Keck/DEIMOS.\\label{tab:sum_obs}} \\tablehead{ \\colhead{Mask ID}    & \\colhead{Field} &  \\colhead{Date (UT)} & \\colhead{Total Exposure}  & \\colhead{$N_z$\\tablenotemark{$\\dagger$}}    & \\colhead{grating}   & \\colhead{central wavelength}   & \\colhead{filter} \\\\ \\colhead{ } & \\colhead{ } & \\colhead{ } & \\colhead{[sec]} & \\colhead{ } & \\colhead{[lines mm$^{-1}$]} & \\colhead{[{\\AA}]} &  \\colhead{ } } \\startdata SDFZD1B & SDF  &   2010 February 13  & $19350$  &  $5$  &  $830$  &  $9000$  &  OG550  \\\\ SDFZD3 & SDF  &   2010 April $14-15$ & $30000$  &  $2$  &  $830$  &  $9000$  &  OG550  \\\\ SDFZD4 & SDF  &   2010 April 15 & $7200$  &  $4$  &  $830$  &  $9000$  &  OG550  \\\\ GNZD1B & GOODS-N  &  2010 February 13, April 14-15 & $18000$  &  $2$  &  $830$ &  $9000$  &  OG550  \\\\ HDF11C & GOODS-N  &  2011 April $1-2$ & $14600$  &  $2$  &  $600$ &  $7500$  &  GG455  \\\\ HDF11D & GOODS-N  &   2011 April 3 & $7200$  &  $2$  &  $830$ &  $8100$  &  OG550 \\enddata \\tablenotetext{$\\dagger$}{% Numbers of observed $z$-dropouts. Some objects were observed on multiple masks (See Table \\ref{tab:sum_followup}). } \\end{deluxetable*}    \\begin{figure*} \\hspace{0.3cm} \\includegraphics[scale=0.75]{./zdrops_2d1dspec.eps} \\caption[] { Spectra of SDF-63544 (left) and SDF-46975 (right). The top panels show the composite two-dimensional spectra, from which the one-dimensional spectra shown in the bottom panels are derived. A prominent emission line is seen at $9683$ {\\AA} for SDF-63544 ($S/N \\simeq 13$), and at $9536$ {\\AA} for SDF-46975 ($S/N \\simeq 14$). } \\label{fig:2d1dspectra} \\end{figure*}    Measuring the fraction of LAEs among Lyman-break Galaxies \\citep[LBGs;][]{giavalisco2002}, the Ly$\\alpha$ fraction, provides complementary information to understand the reionization process \\citep[e.g.,][]{stark2010}. Since LBGs are selected over a broader range of redshifts compared to observations with a narrow-band filter, their number density is less sensitive to cosmic variance. Searching for Ly$\\alpha$ emission from samples of LBGs with available spectra at $4 < z < 6$, \\cite{stark2010b} showed that lower luminosity LBGs have larger Ly$\\alpha$ fractions, and that the fraction increases with redshift \\citep[see also,][]{vanzella2009,stark2010,douglas2010}. However, it is not yet clear if this trend continues $z > 6$. To explore this, we require spectroscopy of dropout candidates at $z \\sim 7$.    The real challenge in estimating the number density of $z > 6$ galaxies and the intensity of ionizing Ly$\\alpha$ photons is the spectroscopic confirmation of these candidates. They are extremely faint, and the only detectable feature is Ly$\\alpha$ emission shifted to near-infrared wavelengths \\citep[e.g.,][]{iye2006,lehnert2010,fontana2010,vanzella2010d}. Despite significant efforts to spectroscopically confirm $z > 6$ candidates, the number of confirmed sources is still very limited. \\cite{fontana2010} reported the detection of one LBG with Ly$\\alpha$ emission at $z=6.97$ in ultra-deep spectroscopy of seven $z$-dropout candidates selected from VLT/Hawk-I imaging of the Great Observatories Origins Deep Survey's southern (GOODS-S) field \\citep{castellano2010}. They found a significant decline in the fraction of LBGs with Ly$\\alpha$ emission between $z \\sim 6$ and $7$, reversing the increasing trend with redshift found at $z<6$. Furthermore, \\cite{vanzella2010d} spectroscopically confirmed two $z$-dropout galaxies at $z \\approx 7.1$ selected from VLT/Hawk-I imaging of the BDF4 field \\citep{castellano2010b}. With the small number of $z$-dropout galaxies with available spectroscopic redshifts, any measure of their Ly$\\alpha$ photon budget or number density will be seriously affected by statistical uncertainties and cosmic variance.    In this study we present results from ultra-deep spectroscopy with Keck/DEIMOS for a sample of $11$ $z$-dropout galaxies. To maximize the spectroscopic success, we designed the photometric survey to identify dropout candidates at the bright end of the UV luminosity function. The aim is to derive the fraction of LBGs with Ly$\\alpha$ emission and to study its evolution to $z \\sim 7$.    In the next section we present the photometric selection of $z$-dropout candidates which are the targets for our spectroscopic observations here. The spectroscopic observations are described in Section \\ref{sec:observations}. This is followed by redshift identification and the measurement of spectroscopic properties in Section \\ref{sec:results}. In Section \\ref{sec:discussion} we discuss the implications for reionization. The conclusions are presented in Section \\ref{sec:conclusion}. Throughout this paper, we use magnitudes in the AB system \\citep{oke1983} and assume a flat universe with ($\\Omega_{\\rm m}$, $\\Omega_{\\rm \\Lambda}$, $h$) $= (0.3, \\, 0.7, \\, 0.7)$. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have presented Keck/DEIMOS spectroscopic observations of $11$ $z$-dropout galaxies found in the SDF and GOODS-N fields. An emission line has been detected at $9500 - 10000${\\AA} in the spectra of three objects, one of which is the previously reported Ly$\\alpha$ emitter IOK-1 at $z = 6.96$. Since all the detected lines are singlet with a large positive weighted skewness, we have concluded that the three objects are Ly$\\alpha$-emitting $z$-dropout galaxies at $z_{\\rm spec} = 7.213$, $6.965$, and $6.844$. Their Ly$\\alpha$ fluxes and rest-frame Ly$\\alpha$ EWs are $2.5 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and $33${\\AA}, $2.8 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and $43${\\AA}, $2.7 \\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and $43${\\AA}, respectively. It should be noted that the $z=7.213$ galaxy was confirmed by observations in two independent DEIMOS runs in $2010$ and $2011$ with three different spectroscopic configurations. This galaxy is detected in HST/WFC3 F140W imaging as well as in Spitzer/IRAC $3.6$ and 4.5 $\\mu$m imaging. Stellar population modeling indicates that this galaxy has a very young age, a stellar mass $\\lesssim 10^9 M_\\odot$, and a high star formation rate of $30 -100 M_\\odot$ yr$^{-1}$, with strong nebular emission contributing to the fluxes in the IRAC bands.    We then have measured the fraction of Ly$\\alpha$-emitting galaxies at $z \\sim 7$. To reduce statistical uncertainties and possible effects of field-to-field variance, we have combined our results with previous $z$-dropout spectroscopic studies \\citep{fontana2010,vanzella2010d} including very recent ones \\citep{schenker2011,pentericci2011}. We have obtained the $z \\sim 7$ Ly$\\alpha$ fraction of $X^{25}_{{\\rm Ly}\\alpha} = 17 \\pm 10${\\%}, and $X^{55}_{{\\rm Ly}\\alpha} = 4 \\pm 4${\\%} for UV-bright galaxies ($-21.75 < M_{\\rm UV} < -20.25$, or $M_{\\rm UV} \\simeq -21$), and $X^{25}_{{\\rm Ly}\\alpha} = 24 \\pm 12${\\%}, and $X^{55}_{{\\rm Ly}\\alpha} = 16 \\pm 9${\\%} for UV-faint galaxies ($-20.25 < M_{\\rm UV} < -18.75$, or $M_{\\rm UV} \\simeq -19.5$). These low values indicate that the fraction of Ly$\\alpha$-emitting galaxies drops from $z \\sim 6$ to $7$, in contrast to the reported increasing trend from $z \\sim 4$ to $6$. We have also found that $X^{25}_{{\\rm Ly}\\alpha}$ drops more strongly in UV-faint galaxies than in UV-bright galaxies. These findings would suggest that the neutral fraction of the IGM significantly increases from $z \\sim 6$ to $7$, and that the increase is stronger around galaxies with fainter UV luminosities, which is consistent with inside-out reionization models where reionization proceeds from high- to low-density environments."
    },
    "1107/1107.0717.txt": {
        "abstract": "We analyze the recently released CoGeNT data with a focus on their time-dependent properties.  Using a variety of statistical tests, we confirm the presence of modulation in the data, and find a significant component at high ($E_{ee} \\gtap 1.5$ keVee) energies. We find that standard elastic WIMPs in a Maxwellian halo do not provide a good description of the modulation. We consider the possibility of non-standard halos, using halo independent techniques, and find a good agreement with the DAMA modulation for Q$_{\\text{Na}} \\approx 0.3$, but disfavoring interpretations with Q$_{\\text{Na}}$ = 0.5. The same techniques indicate that CDMS-Ge should see an $O(1)$ modulation, and XENON100 should have seen 10-30 events (based upon the modulation in the 1.5-3.1 keVee range), unless L$_{\\text{eff}}$ is smaller than recent measurements. Models such as inelastic dark matter provide a good fit to the modulation, but not the spectrum. We note that tensions with XENON could be alleviated in such models if the peak is dominantly in April, when XENON data are not available due to noise. ",
        "introduction": "\\label{sec:intro}  The CoGeNT Collaboration has recently published results from the first fifteen months of data taking~\\cite{Aalseth:2010vx, Aalseth:2011wp}.  Since their first data release more than a year ago, they continue to observe an \u00f5unexplained excess in the spectrum of nuclear recoil scattering rate and now claim an annual modulation of 2.8$\\sigma$\\cite{Aalseth:2011wp}.  This preliminary evidence for a modulation is an important step towards determining the nature of CoGeNT's unexplained spectrum and has been claimed to be evidence for a $\\sim7$ GeV dark matter~\\cite{Hooper:2011hd}.  In this work, we present a comprehensive statistical analysis of the CoGeNT data and show that the modulation spectrum is hard to achieve with a conventional light elastic WIMP in a standard Maxwellian halo.  Direct detection experiments such as CoGeNT search for the scattering of dark matter off nuclei in ground-based detectors.  The spectrum of nuclear recoil energies depends on the mass and scattering cross section of the dark matter, as well as its velocity distribution in the Galaxy.  One of the most distinctive features of a dark matter signal is that it should modulate annually due to the motion of the Earth's rotation about the Sun~\\cite{Drukier:1986tm}.  In particular, the flux of dark matter as observed in the lab frame is larger in the summer, when the Earth is moving in the same direction as the Sun, than in the winter, when the Earth's motion is against that of the Sun~\\cite{Lewin:1995rx}.  For a Maxwell-Boltzmann velocity distribution, the flux peaks 152 days into the year.  Observing an annual modulation in a potential signal is a crucial step in confirming its origin as dark matter.  Direct detection experiments face the challenge of distinguishing dark matter nuclear recoils from a list of potential backgrounds.  In most cases, experiments utilize a combination of ionization, scintillation, or phonon signals to separate out nuclear recoils~\\cite{Gaitskell:2004gd}.  But the possibility of contamination in the nuclear recoil band remains, for instance due to unaccounted for radioactive decays.  If the signal in the nuclear recoil band modulates with the period and phase expected for a dark matter signal, however, it would be a strong indication of the nature of the interaction producing the signal.  To date, only the DAMA~\\cite{Bernabei:2008yi,Bernabei:2010mq} and CoGeNT~\\cite{Aalseth:2010vx, Aalseth:2011wp} experiments have claimed an annual modulation signal.  The DAMA experiment, which uses target crystals of NaI(Tl), claims an $8.9\\sigma$ modulation with period $0.999\\pm0.002$ years and peaking at $146\\pm7$ days.  The CoGeNT experiment, which uses a Ge target, has recently claimed an annual modulation signal with their first fifteen months of data.  For energies ranging from 0.5-3.0~keVee,\\footnote{The notation keVee refers to the ``electron equivalent energy in keV'', which is defined as the reconstructed recoil energy under the assumption that it is carried by an electron. For nuclear recoils, only part of the recoil energy is visible in the detector---an effect that has to be corrected for by dividing the visible energy by a quenching factor---so that the energy threshold for nuclear recoils is higher than that for electron recoils. When referring to a nuclear recoil energy, we will use the notation ``keVnr''.} they observe a maximal modulation with best-fit modulation fraction of $16.6\\pm3.8\\%$, period $347\\pm29$ days, and minimum at Oct. 16$\\pm 12$ days.  In this energy range, the significance is $2.8\\sigma$.  When the CoGeNT data are fit with a dark matter signal in addition to a constant and exponential background, the best-fit dark matter mass is roughly 7-8 GeV with $\\sigma\\sim 10^{-4}$ pb, which is close to the region of parameter space that is consistent with DAMA~\\cite{Aalseth:2010vx, Hooper:2011hd, Arina:2011si}.  The light dark matter interpretation of CoGeNT and DAMA has been challenged by null results from other direct detection experiments, such as XENON100~\\cite{Aprile:2011hi}, XENON10~\\cite{Angle:2007uj,Angle:2011th}, Simple~\\cite{Felizardo:2010mi,Felizardo:2011uw}, and CDMS~\\cite{Ahmed:2009zw,Ahmed:2010wy}.  The compatibility of the XENON and CoGeNT results has been discussed in greater detail in~\\cite{Collar:2011wq,Collar:2010ht,Collar:2010nx,Collar:2010gd,Collar:2010gg}.  Reconciling CDMS and CoGeNT is more challenging because the two use the same target material and CDMS has reported an event rate significantly below that of CoGeNT in the same energy range; however, it has been claimed that errors in the energy calibration can potentially cause the discrepancy~\\cite{Collar:2011kf}. As we shall see, our conclusions on the modulation at CoGeNT will not depend strongly on these details.  This work presents a detailed statistical analysis of the CoGeNT results using the publicly available data~\\cite{juanprivate}. % going beyond the recently presented analyses~\\cite{Frandsen:2011ts,Schwetz:2011xm} which use only the 2 bin data shown in \\cite{Aalseth:2011wp}. Section~\\ref{sec:techniques} introduces the statistical tests that will be used throughout the paper.  Section~\\ref{sec:modulation} presents a model-independent analysis of the modulated and unmodulated rate, period, and phase. %We show that the modulation is most significant in the energy region between 1.5--3.0~keVee; % the best-fit modulation amplitude is larger at low energies (0.5--0.9~keVee), but has a larger error bar.  In the region of greatest significance, the period of the modulation is approximately one year and the phase is $t_0 \\sim 106$ days, although neither the period nor the phase can be well-constrained in energy regions with lower significance. Section~\\ref{sec:darkmatter} discusses the implications of these results for dark matter  and shows that the hypothesis of a light, elastically-scattering WIMP (Weakly Interacting Massive Particle) is strained. In addition, the consistency of CoGeNT with CDMS-Ge, CDMS-Si, XENON100 and DAMA is presented using an analysis that is independent of astrophysical uncertainties. % We use techniques to compare to other experiments independent of the halo model and find that for spin-independent scattering both CDMS-Ge and CDMS-Si indicate that whatever signal is present at high energies should be $\\sim 100\\%$ modulated. The significant modulation at energies above 1.5~keVee would predict large rates at XENON100 unless the $L_{\\rm eff}$ is smaller than recent measurements. At the same time, we find that such a mapping gives remarkable agreement between the rates of modulation at CoGeNT and DAMA, while disfavoring earlier approaches with large $Q_{\\text{Na}}$.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%The search for dark matter is a central element of modern astrophysics, modern cosmology and particle physics. The discovery of particle dark matter is of such importance that any claim must be corroborated by another experiment, and within a single experiment, before it can be believed. We have studied the event rate from the CoGeNT experiment over the first 458 days of running.  CoGeNT continues to see an unexplained excess at low energies and now claims evidence for a modulating spectrum.  The presence of modulation makes the possibility of a WIMP explanation more compelling, but also requires a serious and thorough discussion of its implications for theory and other direct detection experiments.  In this work, we confirm the original findings of \\cite{Aalseth:2011wp} of a modulation in the low-energy (0.5-3.1 keVee) data of CoGeNT, with a significance of  99\\% or 2.6$\\sigma$ (99.7\\% or 3$\\sigma$) for a Maxwell-Boltzmann (best-fit) phase.  However, the details of the modulation make it somewhat confusing. The significance for modulation in the low energy range (0.5-1.5 keVee), where the signal from a light WIMP would be present, is only 87\\% (90\\%) for a Maxwell-Boltzmann (best-fit) phase. In contrast, the significance is 97.7\\% (98.3\\%)  in the higher energy range (1.5-3.1 keVee).  The absence of significant modulation below $\\sim 1.5$ keVee is not troubling on its own, because modulation fractions can be larger at higher energies, and cosmogenic backgrounds at 0.9-1.5 keVee are significant. Nonetheless, one cannot claim that the modulation in the low-energy regime gives evidence for a model while simultaneously disregarding the much more significant modulation in the high-energy regime $-$ i.e., if a model does not explain a significant part of the high energy modulation, it cannot be claimed to explain the modulation. The significance of modulation in the 1.5-3.1 keVee region is confirmed by a variety of statistical techniques, and is an important contributor to the total modulation significance in the 0.5-3.1 keVee energy range.  Many elastic dark matter models with Maxwellian halos do not give any significant modulation above 1.5 keVee for masses below $\\sim 7-9$ GeV. We have performed a broad study of elastic WIMP scenarios with Maxwellian and NFW-consistent halos and have found no models that give modulation at high energies, while not exceeding the {\\em unmodulated} count rate at lower energies.  Halo models with debris flows \\cite{Lisanti:2011as} or streams might allow for this, however.  Our best-fit WIMP point with a Maxwell-Boltzmann halo does not significantly improve the fit compared to a constant background.  As a result, we believe that attempts to understand the CoGeNT modulation with a Maxwellian halo are likely to be unreliable and lead to erroneous conclusions.  To that end, we have attempted to employ techniques that are independent of the halo model \\cite{Fox:2010bz} when comparing the CoGeNT results to other experiments.  We find that a direct comparison to CDMS-Ge allows the modulation, but predicts that a significant modulation should appear at CDMS as well. Note that small errors in energy, while important for interpreting a rapidly falling background, should not affect a general modulation analysis such as this one.  Other experiments also provide insight on the modulation in the 1.5-3.1 keVee range.  In particular, the absence of a signal at CDMS-Si suggests that any signal in this range should be $\\sim$ 100\\% modulated.  A study of XENON100 is intriguing, because the modulation appears in a range that has been more directly calibrated with measurements of L$_{\\text{eff}}$. We find that any significant modulation in the 1.5-3.1 range should have shown up at the 10-30+ event level at XENON100, unless the value of L$_{\\text{eff}}$ is significantly lower than what has been found by \\cite{Plante:2011hw}.  A halo model-independent comparison to DAMA shows generally {\\em good} agreement between the event rates for a light WIMP and Q$_{\\text{Na}} \\approx 0.3$.  However, the modulation at CoGeNT is in conflict with DAMA if one assumes Q$_{\\text{Na}} = 0.5$, disfavoring previous interpretations utilizing a Maxwellian halo.  Inelastic dark matter models can provide a good fit to the modulation, but not the exponentially falling rate at low energies. iDM in the presence of stream(s) or debris flows could lead to the highly-modulated signal that is observed. If the signal is only present in April, it could explain why XENON100 did not observe a significant rate, as data taking in that period was limited by excess noise.  In summary, the modulation at CoGeNT is an intriguing piece of the puzzle, but raises as many questions through the high-energy modulation as it answers.  Future data from CoGeNT is imperative for understanding the presence or absence of modulation at low and high energies and thus, shedding light on whether--and what type of--dark matter could give rise to it.  \\noindent {\\bf Note Added} At the time of the completion of the manuscript, the authors became aware of related papers~\\cite{Schwetz:2011xm, Farina:2011pw}, which explore implications of the CoGeNT data for particle and astrophysical models.  The authors of \\cite{Schwetz:2011xm} carried out an analysis on the two bin data presented in \\cite{Aalseth:2011wp}, while \\cite{Farina:2011pw} use the full data set analysed in this work.  Our conclusions are in qualitative agreement.  After this manuscript had been submitted to the e-print arXiv, the CoGeNT collaboration released new preliminary results which indicate that part of the low-energy excess events may be due to surface backgrounds~\\cite{Collar-TAUP}. If confirmed, this would imply that the parameter region favored by the CoGeNT event spectrum would become larger and move to lower cross sections~\\cite{Collar-TAUP, Kopp:2011yr, Kelso:2011gd}.   \\vskip 0.2in \\noindent {\\bf Acknowledgements} We are especially grateful to Juan Collar for making the CoGeNT data publicly available and for his support in explaining many aspects of it during our analysis. We also wish to thank Rouven Essig, Roni Harnik, Graham Kribs, Rafael Lang, Josh Ruderman, Tracy Slatyer, Natalia Toro, Michael Witherell and Itay Yavin for useful discussions.  ML acknowledges support from the Simons Postdoctoral Fellowship and the LHC Theory Initiative. NW is supported by NSF grant \\#0947827, as well as support from the Amborse Monell Foundation. This research was supported in part by the National Science Foundation under Grant No. NSF PHY05-51164. Fermilab is operated by Fermi Research Alliance, LLC, under Contract DE-AC02-07CH11359 with the United States Department of Energy.  %%%%%%%%%%%%%%%%%%  \\begin{appendix}"
    },
    "1107/1107.1181_arXiv.txt": {
        "abstract": "We conduct a descriptive analysis of the multipolar structure of gravitational-radiation waveforms from equal-mass aligned-spin mergers, following an approach first presented in the complementary context of nonspinning black holes of varying mass ratio [Baker \\etal Phys. Rev. D \\textbf{78}, 044046 (2008)]. We find that, as with the nonspinning mergers, the dominant waveform mode phases evolve together in lock-step through inspiral and merger, supporting the previous waveform description in terms of an adiabatically rigid rotator driving gravitational-wave emission -- an \\emph{implicit rotating source} (IRS). We further apply the late-time merger-ringdown model for the \\emph{rotational frequency} introduced in Baker \\etal (2008), along with an improved amplitude model appropriate for the dominant $(2,\\pm 2)$ modes. This provides a quantitative description of the merger-ringdown waveforms, and suggests that the major features of these waveforms can be described with reference only to the intrinsic parameters associated with the state of the final black hole formed in the merger. We provide an explicit model for the merger-ringdown radiation, and demonstrate that this model agrees to fitting factors better than 95\\% with the original numerical waveforms for system masses above $\\sim 150 \\MSun$. This model may be directly applicable to gravitational-wave detection of intermediate-mass black-hole mergers. ",
        "introduction": "\\label{sec:intro}  Black-hole-binary mergers are a key target of ground-based and space-based gravitational-wave observations.  The strongest radiation is produced just as the two black holes join to become one, and can only be fully understood through explicit numerical simulations. Since the first stable evolutions of black-hole-binary mergers \\cite{Pretorius:2005gq,Pretorius:2006tp,Baker:2005vv,Campanelli:2005dd}, and after it was established that the gravitational waveforms from these evolutions were universal, and consistent across codes and methodologies \\cite{Baker:2006yw,Baker:2007fb,Hannam:2009hh}, researchers have turned their attention to how the results of numerical relativity can most usefully be supplied to the gravitational-wave data-analysis community.  After studying the equal-mass nonspinning case, researchers have had to address the complexity problem of more generic systems. Even allowing for simple scaling by total mass $M = M_1 + M_2$, and assuming zero eccentricity, such systems span a seven-dimensional parameter space: $\\{\\eta, \\vec{S_1}, \\vec{S_2}\\}$, where $\\eta = M_1 M_2/M^2$ is the \\emph{symmetric mass ratio} of the binary, and $\\vec{S_i}$ is the spin angular momentum vector of hole $i$.  Early surveys of the waveform parameter space have restricted themselves to the $\\eta$-dependence of nonspinning systems. In \\cite{Baker:2008mj}, the authors investigated the multipole structure of merger waveforms from such systems, noting that the strongest subdominant modes shared many characteristics with the dominant quadrupole, and that they could be collectively described by an \\emph{implicit rotating source} model of the binary. The authors used this observation to construct a multi-mode gravitational-wave template family for such binary systems, as an alternative to more usual effective-one-body (EOB) templates \\cite{Buonanno:2007pf,Buonanno:2009qa}.  While we may assume that $\\eta$ and $|\\vec{S_i}|$ remain essentially constant throughout inspiral and merger, the spin directions generally evolve, so a useful parametrization of the system should take care to distinguish components of the spin-direction space with physically distinct effects on the waveforms \\cite{Galley:2010rc}.  An obvious cut in parameter space to consider is that of spins aligned (or anti-aligned) with the orbital angular momentum. These systems will not precess, but exhibit observationally significant spin-orbit effects, distinguishing them from nonspinning binaries in their dynamics and resulting waveforms \\cite{Campanelli:2006uy}. High-accuracy waveforms from such evolutions have been produced and studied by several groups \\cite{Shoemaker:2008pe,Reisswig:2009vc,Chu:2009md,Mosta:2009rr}. Such systems have been partially characterized by \\cite{Hannam:2010ec}, using a variant of the frequency model from \\cite{Baker:2008mj}. The frequency-domain phenomenological templates of Ajith \\etal have been extended to cover both mass ratio and total aligned spin \\cite{Ajith:2009bn,Santamaria:2010yb}, at least for the dominant modes, and attempts have been made to extend these to more generic systems \\cite{Sturani:2010ju,Sturani:2010yv}.  A key result from our investigation of the dominant modes of nonspinning unequal-mass binary waveforms \\cite{Baker:2008mj} was that these modes had phases that evolved together in lock-step through inspiral and merger. This agreement was especially impressive for the $\\ell=m$ modes, leading to the development of a heuristic picture of the binary system as a  rigid rotator (at least in the adiabatic limit) driving gravitational-wave emission. We dubbed this the \\emph{implicit rotating source} (IRS) picture.  A secondary result of this picture was the possibility of developing a simple model for the time-development of the dominant and leading subdominant modal frequencies $\\omega_{\\ell m}$ in terms of a single \\emph{rotational frequency} $\\Omega(t)$: \\[ \\omega_{\\ell m} = m \\Omega(t). \\] We also presented a simple model for the corresponding mode amplitudes, leading to the possibility of a new approach to time-based gravitational-waveform templates. In fact, we developed such a template proposal, the IRS-EOB templates, as an alternative to the effective-one-body templates of \\cite{Buonanno:2007pf,Buonanno:2009qa}, which terminate the signals by matching to a superposition of quasinormal-mode (QNM) frequencies.  In this paper, we look at the dominant waveform modes from some aligned-spin systems, and ask the following general questions: Does the general IRS picture still hold? How do the features of aligned-spin mergers compare with those of nonspinning mergers? Can we quantify the main features of the merger ringdown with a simple analytic model?  The rest of this paper is laid out as follows. In Sec.~\\ref{sec:sims}, we introduce the binary systems studied and the numerical methods used to simulate them. In Sec.~\\ref{sec:radiation}, we present results for the final black-hole states and an IRS descriptive characterization of the waveforms from our numerical evolutions. In Sec.~\\ref{sec:model}, we analyze the late portions of these waveforms in more detail, and apply the analytic modeling approach of \\cite{Baker:2008mj} to the dominant-mode frequencies and (with improvements) to the amplitude model, concluding with an explicit parametrization approximating the $(2,2)$ results of all our simulations. In Sec.~\\ref{sec:match}, we investigate the quality of the new models compared to the numerical waveforms in the context of the Advanced LIGO detector. We conclude with some discussion in Sec.~\\ref{sec:discuss}. Some extra detail on the convergence of the numerical simulations is given in the Appendix. ",
        "conclusions": "\\label{sec:discuss}  In this paper, we have investigated the mode-decomposed gravitational waveforms resulting from the merger of aligned-spin black-hole binaries. Our primary purpose was to establish how well the implicit rotating source (IRS) picture of the binary as a GW source -- first suggested in \\cite{Baker:2008mj} in the context of nonspinning holes of comparable mass ratios -- holds in this different branch of parameter space.  Based on these investigations, we note that the modal structure of aligned-spin mergers is like that of the equal-mass nonspinning configuration, dominated by  the $(2,\\pm 2)$, $(3,\\pm 2)$, and $(4,\\pm 4)$ modes.  These modes still display IRS-type behavior, featuring common rotational phase evolution with little offset through late-inspiral, merger, and into ringdown. The peak modal amplitudes are similar to those for the nonspinning case, though the duration of the peak region (which was roughly independent of mass ratio) is extended for aligned spins (and shortened for anti-aligned spins). A similar timescale dependence is seen in the rise to peak frequency.  In applying our late-merger frequency model (with a slightly modified parametrization) to these new cases, we have found that the model still performs well for the dominant modes. We enhance the original mode amplitude model of \\cite{Baker:2008mj} to achieve improved behavior, at least for the leading $(2,\\pm 2)$ modes; however it yields up to $\\sim 10\\%$ mismatches at and before peak for the next most important modes.  For the $(2,\\pm 2)$ modes at least, we have attempted to constrain all parameters explicitly with reference only to the state of the final black hole (i.e., its dominant quasinormal-mode complex frequency). With these constraints, we have reduced the additional free parameters to just three: $t_0$, the time of maximum chirp-rate, $\\varphi_0$, the phase offset, and $A_0$, the mode's amplitude scale. This description provides an approximate fit to the late part of the waveforms for all our simulations, including equal-mass spinning cases, and the 4:1 nonspinning case.  Moving back in time to earlier points before the merger, the quality of this fit degrades and other physical details of the premerger binary become more significant.  We see evidence of this when comparing the \\texttt{X4\\_00} (4:1 nonspinning) and the \\texttt{X1\\_DD} (equal-mass down-down spins) configurations, which result in the same final spin.  We have quantified the quality of this approximation by calculating Advanced LIGO fitting factors. For system masses of $\\gtrsim 150 \\MSun$, we have found mismatches of $\\lesssim 5\\%$ between the full numerical-relativity waveform and the $(2,\\pm 2)$-mode-only model waveforms.  Our results suggest that an approach to gravitational-wave observation templates with parameters tied first to the structure of the final black hole may be useful for Advanced LIGO observations of intermediate-mass mergers. These would be an alternative to the time-domain effective-one-body (EOB) templates of \\cite{Pan:2007nw,Pan:2009wj,Pan:2010hz}, and the frequency-domain ``phenomenological'' templates of \\cite{Ajith:2009bn,Santamaria:2010yb}, similar to ringdown searches currently being applied to LIGO data \\cite{LSC-Virgo_GWDA,Goggin:2008dz,Cadonati:2009jg,Sengupta:2009nm}.  In future work we plan to investigate the quality of this model, or its extensions, for a broader variety of mergers, including precessing and eccentric configurations.  In this explicit modeling, we have focused on the dominant $(2,2)$ mode. A similar model incorporating full multi-mode information can also be applied to complete inspiral-merger-ringdown waveform templates, as was done for nonspinning systems in \\cite{Baker:2008mj}. Future work on this topic will focus on improving and extending the amplitude model to cover multiple significant modes of the merging binary."
    },
    "1107/1107.3395.txt": {
        "abstract": "{Rapidly rotating stars show short-period oscillations in magnetic activity and polar appearance of starspots.}{The aim of this paper is to study large-scale shallow water waves in the tachoclines of rapidly rotating stars and their connection to the periodicity and the formation of starspots at high latitudes.}{Shallow-water magnetohydrodynamic equations were used to study the dynamics of large-scale waves at the rapidly rotating stellar tachoclines in the presence of toroidal magnetic field. Dispersion relations and latitudinal distribution of wave modes were derived.}{We found that low-frequency magnetic Rossby waves tend to be located at poles, but high-frequency magnetic Poincar\\'e waves are concentrated near the equator in rapidly rotating stars. These results have important implications for the evolution of the stellar wind in young Sun-like stars.}{Unstable magnetic Rossby waves may lead to the local enhancement of magnetic flux at high latitudes of tachoclines in rapidly rotating stars. The enhanced magnetic flux may rise upwards owing to the magnetic buoyancy in the form of tubes and appear as starspots at polar regions. Magnetic Rossby waves may also cause observed short-term periodicity in the stellar magnetic activity. } ",
        "introduction": "Rotation and convection lead to stellar dynamo activity, which results in the strong concentration of magnetic fields on the surface. These magnetic features are called {\\it starspots}, and they represent the sunspot analog for other stars. Interested readers may find detailed information about starspots in two recent comprehensive reviews: Berdyugina (\\cite{Berdyugina2005}) and Strassmeier (\\cite{Strassmeier2009}). Hall (\\cite{Hall1972}) explicitly postulated the starspot model as explaining the wave-like features in the light curves of rotating RS CVn stars. Doppler imaging technique showed spots on the RS CVn star HR 1099, and they exhibit a similarity in shape and location to X-ray images of solar coronal holes (Vogt and Penrod, \\cite{Vogt1983}). Later, Strassmeier (\\cite{Strassmeier1990}) observed a cool polar spot on the rapidly rotating G5 subgiant in the single-lined RS CVn binary HD 26337. The spot was rather long-lived, persisting for at least the nine years of observations.  Zeeman-Doppler imaging of the circularly polarized spectrum has led to the first magnetic maps of the active K1 subgiant in the RS CVn binary system HR 1099 = V711 Tau (K1IV+G5V) other than the Sun, which also showed the polar location of starspots (Donati et al. \\cite{Donati1992}). It was generally accepted that starspots on magnetically active, cool stars mainly appear near the poles. Sch\\\"ussler and Solanki (\\cite{Schussler1992}) suggested that the preference for high latitudes is due to the rapid rotation to these stars, which in turn leads to a dominance of the Coriolis force over the buoyancy force in the dynamics of magnetic flux tubes. %As a consequence, flux tubes erupting from the deep parts of the stellar convection zone follow a path nearly parallel to the axis of rotation and %thus necessarily surface at high latitudes, unless their initial field strength exceeds a critical value for which buoyancy becomes dominant again. Later, Sch\\\"ussler et al. (\\cite{Schussler1996}) carried out simulations on solar-type stars in an effort to predict the emergence latitude of magnetic flux tubes, after evolution from a toroidal magnetic field located at the base of the convection zone. They carried out simulations for stars of different ages and equatorial rotation velocities. The results predict high-latitude spots since, in a fast rotator, Coriolis forces acting on a rising magnetic flux tube cause movement parallel to the rotation axis.  On the other hand, Barnes et al. (\\cite{Barnes1998}) studied two rapidly rotating G dwarfs in the $\\alpha$ Persei cluster, from spectra taken at the William Herschel Telescope. The image reconstructions demonstrate that both these stars have cool high-latitude regions or polar crowns and low-latitude features, contradicting the suggestion that only high-latitude spots should be present. The magnetic topology of AB Dor (young K0 dwarf) is also found to be very complex, with at least 12 different radial field regions of opposite polarities located all around the star (Donati et al. \\cite{Donati1999, Donati2003}). Significant azimuthal field fluxes are also detected in the form of one negative polarity region close to the equator, a series of positive polarity patches at intermediate latitudes and an almost complete ring of negative polarity encircling the rotational pole at high latitudes. Starspots on the young solar analog EK Dra (rotation period of 2.68 days) are observed to be formed everywhere, from the low latitudes up to polar regions (Strassmeier 2009). Therefore, the current view is that starspots appear at all latitudes in rapidly rotating stars.  The high-latitude sunspots could change the structure of coronal magnetic field as well as the formation of prominences, which may lead to the predominance of solar flares and coronal mass ejections (CMEs) near the polar regions. CMEs are believed to have tremendous influence on the atmosphere of exoplanets, leading to their erosion (Khodachenko et al. \\cite{Khodachenko2007}). The polar CMEs may propagate significantly out off the ecliptic plane, which may cause more convenient conditions of planetary evolution; therefore, the detailed study of starspot distribution in rapidly rotating young solar analogs is very important for planetary evolution as well.  Observation of starspots at all latitudes complicates any explanation of their appearance. There are at least two suggested mechanisms of starspot formation in rapidly rotating stars. The first mechanism, which we already mentioned above, is that the magnetic tubes tend to move to the polar regions owing the Coriolis force acting during the rising across the convection zone (Sch\\\"ussler and Solanki \\cite{Schussler1992}; Sch\\\"ussler et al. \\cite{Schussler1996}; DeLuca et al. \\cite{DeLuca1997}; Granzer et al. \\cite{Granzer2000}; Holzwarth et al. \\cite{Holzwarth2006}, Isik et al. \\cite{Isik2007}). The second mechanism is the transport of photospheric magnetic flux towards the polar regions by meridional advection  (Schrijver and Title \\cite{Schrijver2001}, Mackay et al. \\cite{Mackay2004}). The first mechanism may explain the starspot appearance at high latitudes, but it may fail to explain the equatorial spots because the magnetic tubes should all tend to higher latitudes in rapid rotators. On the other hand, the second mechanism may explain the distribution of spots at all latitudes, but it requires a strong meridional flow in order to carry the magnetic flux towards the poles.  Another interesting observational indication is the cyclic behavior of stellar activity with shorter periods than the solar cycle ($\\sim$ 11 yr). Massi et al. (\\cite{massi1998}, \\cite{massi2005}) report short-term activity cycles on UX Arietis, which is a binary system formed by a spotted K0 IV primary and a G5 V secondary with an orbital period of 6.44 days. Short-period activity cycles have been observed on $\\tau$ Bootis (Donati et al. \\cite{Donati2008}). Ol\\'ah et al. (\\cite{Olah2009}) analyzed 20 active stars and came to the conclusion that stellar activity cycles are generally multiple and variable. Recently, Lanza et al. (\\cite{Lanza2009}) have observed a short-term cyclic oscillation of the total spotted area on the photosphere of CoRoT-Exo-2a, which is supposedly caused by Rossby waves. The rotation period of CoRoT-Exo-2a is 4.522 $\\pm$ 0.024 days, while the total spotted area shows a cyclic oscillation with a period of 28.9 $\\pm$ 4.3 days. Almost all observed periodicities are shorter than the solar cycles and are several times longer than the stellar rotation periods. It should be mentioned, however, that solar activity also undergoes intermediate periodicities in the range of several months (Rieger et al. \\cite{Rieger84}, Carbonell et al. \\cite{Carbonell1990}, Oliver et al. \\cite{oliver1998}) and $\\sim$ 2 years (Sakurai \\cite{sakurai81}, Gigolashvili et al. \\cite{gigolashvili95}, Vecchio et al. \\cite{vecchio10}). These periodicities are explained in terms of magnetic Rossby wave instability in the solar tachocline that is the joint effect of latitudinal differential rotation and toroidal magnetic field (Zaqarashvili et al. \\cite{Zaqarashvili2010a}, \\cite{Zaqarashvili2010b}). The coincidence between the solar short term oscillations and the cyclic behavior of spotted stars indicate that magnetic Rossby waves can be important in rapidly rotating stars.  Here we study the behavior of magnetic shallow water waves in tachoclines of rapidly rotating stars and their role in the latitudinal distribution of starspots. The tachocline is the thin layer located between the external convection envelope and the internal radiative core (in the case of the Sun see Spiegel and Zahn \\cite{Spiegel1992}). Rossby waves are well studied in the geophysical context (Gill \\cite{gill82}); however, the presence of magnetic fields significantly modifies their dynamics (Zaqarashvili et al. \\cite{Zaqarashvili2007}, \\cite{Zaqarashvili2009}, Heng and Spitkovsky \\cite{Heng2009}, Mathis and de Brye \\cite{mathis2011}). We have used the shallow water magnetohydrodynamic (MHD) equations (Gilman \\cite{Gilman2000}) to study the latitudinal structure of magnetic shallow water waves in tachoclines of rapidly rotating stars. We first write the governing equations of shallow water MHD in the rotating frame. Then, we derive the dispersion relations and latitudinal structure of magnetic shallow water waves in different configurations of toroidal magnetic field, which is followed by the discussion of the connection between the shallow water waves and the starspots at stellar photospheres.  Finally we discuss the implications of our findings in the context of the evolution of young stellar winds of Sun-like stars. There is observational evidence that zero-age-main-sequence solar-type stars rotate more than ten times faster than the present Sun (e.g., G\\\"udel et al. \\cite{Gudel1997}, Ribas et al. \\cite{Ribas2005}, Lammer et al. \\cite{Lammer2011}). As a consequence, young solar-type GV stars, including the early Sun, should have vigorous magnetic dynamos and correspondingly strong high-energy emissions. From the study of solar-type stars with different ages, Skumanich (\\cite{Skumanich1972}), Simon et al. (\\cite{Simon1985}) and others have shown that these young stars lose angular momentum with time via magnetized stellar winds, thus leading to a secular increase in their rotation period (Durney \\cite{Durney1972}). This rotation slowdown is well fitted by a power law roughly proportional to $t^{1/2}$, where $t$ is the  stellar age in units of Gyr (e.g., Skumanich \\cite{Skumanich1972}; Soderblom \\cite{Soderblom1982}; Ayres \\cite{Ayres1997}). As a response to slower rotation the solar dynamo strength diminishes with time, causing the Sun's high-energy emissions also to undergo a significant decrease (Ribas et al. \\cite{Ribas2005}). The fundamental question of whether the young fast-rotating Sun and similar stars have experienced some periods of stronger magnetic activity has a major impact on the evolution and plasma interaction of planetary atmospheres in the Solar System and beyond. ",
        "conclusions": "Large-scale magnetic Rossby waves may play a significant role in the dynamics of rapidly rotating stars. Here we have studied the shallow-water MHD waves in the tachoclines of rapid rotators. We considered a rotating spherical coordinate system and a toroidal magnetic field in the tachocline. We derived the dispersion relations for high (magnetic Poincar\\'e waves) and low (magnetic Rossby waves) frequency spectra separately in the approximation of rapid rotation. We found that magnetic Rossby waves tend to locate mainly at the polar regions, while magnetic Poincar\\'e waves are located at low latitudes. Magnetic Rossby waves may lead to the enhancement of magnetic flux at the polar regions and trigger its eruption towards the surface in the form of magnetic flux tubes. Then the magnetic tubes may form starspots at the polar region. Magnetic Rossby waves may cause the observed intermediate period oscillations in the magnetic activity of rapidly rotating stars, such as young solar proxies. The polar location of prominences during the fast-rotating period of the young star may trigger polar CMEs, which propagate significantly out of the ecliptic plane and completely change the current picture of the solar/stellar plasma interaction with young planetary environments."
    },
    "1107/1107.1238_arXiv.txt": {
        "abstract": "We present axisymmetric, orbit-based models to study the central black hole, stellar mass-to-light ratio, and dark matter halo of NGC~4594 (M104, the Sombrero Galaxy).  For stellar kinematics, we use published high-resolution kinematics of the central region taken with the \\emph{Hubble Space Telescope}, newly obtained Gemini long-slit spectra of the major axis, and integral field kinematics from the SAURON instrument. At large radii, we use globular cluster kinematics to trace the mass profile and apply extra leverage to recovering the dark matter halo parameters.  We find a black hole of mass \\mbh$=(6.6 \\pm 0.4) \\times 10^8 \\, M_{\\odot}$, and determine the stellar $M/L_I=3.4 \\pm 0.05$ (uncertainties are the 68\\% confidence band marginalized over the other parameters). Our best fit dark matter halo is a cored logarithmic model with asymptotic circular speed $V_c=376 \\pm 12 \\text{ km s }^{-1}$ and core radius $r_c= 4.7 \\pm 0.6$ kpc. The fraction of dark to total mass contained within the half-light radius is $0.52$.  Taking the bulge and disk components into account in our calculation of  $\\sigma_e$ puts NGC~4594 squarely on the $M$-$\\sigma$ relation. We also determine that NGC~4594 lies directly on the $M$-$L$ relation. ",
        "introduction": "Most galaxies are thought to host supermassive black holes (SMBHs) at their centers.  The masses of these SMBHs have been observed to correlate with several properties of their host elliptical galaxies and of the classical bulge components of their host disk galaxies.  For example, \\mbhs correlates with galaxy/bulge mass \\citep{dre89,mag98,lao01,mcl02,mar03,har04}, luminosity (the $M$-$L$ relation) \\citep{kor93,kor95,kor01, gul09}, velocity dispersion (the $M$-$\\sigma$ relation) \\citep{fer00,geb00,tre02, gul09}, and globular cluster content \\citep{bur10,har11}.  These and other, similar correlations suggest that galaxy formation and black hole growth are fundamentally linked.  To better understand this interplay, accurate black hole masses are needed.  One challenge that limits the accuracy is the determination of the host galaxy's inclination. Projection effects are difficult to model and cause loss of information, leading to systematic uncertainties. Therefore, SMBHs in galaxies whose inclination is confidently known have the best chance of being accurately and robustly measured.  Another issue that limits the accuracy is the effect a dark matter halo has on the determination of SMBH mass. \\citet{geb09} show that orbit-based models can underestimate black hole mass when dark matter is not considered in the modeling, however \\citet{sch11} find the effect is small when the black hole's sphere of influence is well-resolved.  NGC~4594 (M104, or the Sombrero Galaxy) is a nearly edge-on Sa type spiral with a prominent stellar disk and large, classical bulge \\citep{kor04c}. The shape of this disk indicates that it (and thus the entire galaxy) is inclined at an angle very close to $90^{\\circ}$. Throughout this paper we assume a distance to NGC~4594 of 9.8 Mpc, calculated from surface brightness fluctuations \\citep{ton01}.  Unless otherwise stated, all distance-dependent quantities are scaled to this value. \\citet{ton01} use a value of $H_0=74 \\text{ km s}^{-1}~\\text{Mpc}^{-1}$ in their distance determinations, however we compare our \\mbhs and $L_V$ to \\citet{gul09} who adopt $H_0=70$ in their work. We therefore scale the \\citet{gul09} distances down by $6\\%$.  Black hole mass scales as \\mbh$\\propto D$ and luminosity as $L_V \\propto D^{-2}$; these quantities are adjusted accordingly.  NGC~4594 was one of the first galaxies in which a black hole was detected, and it has a long history of study. \\citeauthor{kor88}(1988, hereafter \\citetalias{kor88}) first found evidence for a massive black hole of \\mbh$=5.4^{+11.8}_{-3.7} \\times 10^8$ $M_{\\odot}$ using only ground-based observations.  With isotropic Jeans models, \\citet{ems94b} measured a black hole of mass \\mbh$=(5.4 \\pm 0.5) \\times 10^8$ $M_{\\odot}$. Later, \\citeauthor{kor96} (1996, hereafter \\citetalias{kor96}) used high-resolution kinematics from the Faint Object Spectrograph (FOS) on the \\emph{Hubble Space Telescope} (\\emph{HST})---the same data set we include---to measure $\\text{log } M_{\\bullet}=8.8 \\pm 0.5 \\, M_{\\odot}$. This corresponds to a mass of $5.8^{+12.4}_{-4.0} \\times 10^8 \\, M_{\\odot}$. With isotropic models, \\citet{mag98} obtained a value of \\mbh$=6.9^{+0.2}_{-0.1} \\times 10^8$ $M_{\\odot}$. These values for \\mbhs all lie towards the high mass end of the $M$-$\\sigma$ and $M$-$L$ relations.  Massive SMBH measurements are frequently being revised, and we expect the confidently known inclination of NGC~4594 to lead to one of the more secure measurements of a high mass SMBH.  We present new Gemini spectroscopy of the major axis, as well as SAURON integral field kinematics covering the central region of the galaxy. We also use high-resolution \\emph{HST}/FOS kinematics of the nucleus and kinematics derived from globular clusters at large radii.  We combine these kinematic datasets with  \\emph{HST} and ground-based photometry to run axisymmetric orbit-based models. These models allow us to measure the black hole mass, stellar mass-to-light ratio, and dark matter halo of NGC~4594.  In addition, we also recover information about the internal orbit structure of the galaxy. ",
        "conclusions": "\\subsection{Black Hole-Bulge Correlations}  We discuss the position of NGC~4594 on the $M$-$\\sigma$ and $M$-$L$ relations, as defined by \\citeauthor{gul09} (2009b, hereafter \\citetalias{gul09}) and compare our values of the correlation parameters to previous measurements.  We calculate the effective velocity dispersion $\\sigma_e$ similar to \\citetalias{gul09}    \\begin{equation} \\sigma_e^2 \\equiv \\frac{{\\int}^{R_e}_{R_{\\mathrm{inf}}} (V^2(R) + \\sigma^2(R)) I(R) dR}{{\\int}^{R_e}_{R_{\\mathrm{inf}}} I(R) dR} \\label{sigmae} \\end{equation}  \\ni where $V(R)$ is the rotational velocity and $I(R)$ is the surface brightness profile. This makes $\\sigma_e$ essentially the surface-brightness-weighted second moment. Instead of integrating from the center of the galaxy ($R=0$) as \\citetalias{gul09} did, we integrate from $R > R_{\\mathrm{inf}}$ to ensure that we do not bias $\\sigma_e$ with the high dispersion near the black hole.  Our outermost kinematic data point is at $R=45\\arcsec$, thus we have a gap in kinematic coverage between $45\\arcsec < R < R_e=114\\arcsec$.  The velocity dispersion near the end of our long-slit data is dropping sharply, however $V$ may still contribute to the integral for $R>45\\arcsec$.  To investigate this, we use the gas rotation curve presented in \\citet{kor89} which extends well beyond $R_e$.  Truncating the integral at $R=45\\arcsec$ gives $\\sigma_e=292 \\text{ km s}^{-1}$ while using the extended rotation curve yields $\\sigma_e=297 \\text{ km s}^{-1}$.  The problem with this definition of $\\sigma_e$ is that it includes a contribution from the rotation of the disk.  It has been shown that black hole mass does not correlate with disk properties \\citep{kor11} so this is not ideal.  However, to compare with \\citetalias{gul09} we must be consistent in our calculation of $\\sigma_e$.  We therefore quote this value of $\\sigma_e$ when we compare to the $M$-$\\sigma$ relation determined by \\citetalias{gul09}.  As we expect black hole mass to track bulge quantities, disk contribution to $\\sigma_e$ is likely to add a source of intrinsic scatter to spiral galaxies in the $M$-$\\sigma$ relation.  In fact, spiral galaxies are observed to have larger scatter about $M$-$\\sigma$ than ellipticals of similar $\\sigma_e$.  We also discuss some possible alternatives to $\\sigma_e$ where we attempt to remove the disk contribution.  One option is to remove $V^2(R)$ from Equation (\\ref{sigmae}) altogether. Bulges are known to rotate, however, \\citep{kor82} and this will likely underestimate $\\sigma_e$. This crude calculation gives $\\sigma_e=200 \\text{ km s}^{-1}$.  Another option is to assume some degree of bulge rotation a priori.  If we assume NGC~4594 rotates isotropically \\citep{kor82}, then its flattening determines its position on the $V/\\sigma-\\epsilon$ diagram \\citep{bin78}. \\citet{kor82b} shows that the relation  \\begin{equation} \\frac{V}{\\sigma} \\approx \\sqrt{\\frac{\\epsilon}{1-\\epsilon}} \\label{v/sigma} \\end{equation}  \\ni approximates the isotropic rotator line to roughly $1\\%$ accuracy. We use our value of the bulge ellipticity $\\epsilon=0.25$ in Equation (\\ref{v/sigma}) and assume this value of $V/\\sigma$ applies globally to the entire bulge.  We then use our measured dispersion profile $\\sigma(R)$ to determine the bulge velocity $V_{\\mathrm{bulge}}(R)$.  Using these quantities, we determine $\\sigma_e=230 \\text{ km s}^{-1}$.  We compare this to the kinematics listed in \\citet{kor82}.  These data include long-slit spectra taken at a position angle parallel to the major axis, but $30\\arcsec$, $40\\arcsec$, and $50\\arcsec$ above it.  From these data, it is apparent that the bulge $\\sigma$ off the major axis is roughly constant at $\\sim 220 \\text{ km s}^{-1}$.  The rotation velocity rises from 0 to $100 \\text{ km s}^{-1}$ at large radii. We estimate the luminosity-weighted mean $V \\sim 50 \\text{ km s}^{-1}$. Adding this in quadrature to the constant bulge $\\sigma=220 \\text{ km s}^{-1}$ gives $\\sigma_e \\approx 226 \\text{ km s}^{-1}$.  This estimate does not contain any rotation from the disk, and is consistent with our determination of $\\sigma_e=230 \\text{ km s}^{-1}$ obtained by assuming a constant $V/\\sigma$.  \\begin{figure*}[t] \\centerline{\\includegraphics[width=15cm,angle=0]{msigma.eps}} \\figcaption[msigma.eps]{Position of NGC~4594 on the \\citetalias{gul09} $M$-$\\sigma$ and $M$-$L$ relations.  The plot of $M$-$\\sigma$ (left) shows the three ways we calculate $\\sigma_e$ as well as the value from \\citetalias{gul09} (black triangle).  In order of increasing $\\sigma_e$ we plot $\\sigma_e$ with no rotation (red diamond), $\\sigma_e$ assuming a value of $V/\\sigma$ (green square) and $\\sigma_e$ as in \\citetalias{gul09} (blue asterisk).  For the $M$-$L$ relation (right) we plot the \\citetalias{gul09} value (black triangle) along with our measurement (green square). \\label{msigma}} \\end{figure*}   \\begin{figure}[t] \\includegraphics[width=9cm,angle=0]{fracdm.eps} \\figcaption[fracdm.ps]{Fraction $M_{DM}/(M_{\\star}+M_{DM})$ of enclosed mass that is dark matter as a function of radius. \\label{fracdm}} \\end{figure}   Our black hole mass \\mbh$=(6.6 \\pm 0.4) \\times 10^8 \\, M_{\\odot}$ agrees nicely with that of \\citetalias{gul09} which uses the \\citetalias{kor88} value. In fact, when corrected for their various distance determinations, most values of \\mbhs in the literature agree quite well (\\citetalias{kor88}, \\citealt{ems94b}, \\citetalias{kor96}, \\citealt{mag98}) despite the many modeling techniques and datasets used.  This is likely due to the high degree of isotropy as evidenced in Figure \\ref{vr_vt}, deduced from the $V/\\sigma-\\epsilon$ diagram \\citep{kor82}, and noted in \\citetalias{kor88}.    Figure \\ref{msigma} plots the position of NGC~4594 on the \\citetalias{gul09} $M$-$\\sigma$ and $M$-$L$ relations.  We plot each determination of $\\sigma_e$ in the left-hand panel.  Straightforward application of Equation (\\ref{v/sigma}) leads to a value of $\\sigma_e$ that falls directly on the \\citetalias{gul09} $M$-$\\sigma$ line (blue asterisk).  Next closest is the method of calculating $\\sigma_e$ by assuming a value of $V/\\sigma$ (green square).  This point lies $0.44$ dex above the \\citetalias{gul09} line, however this is still within the estimated scatter.  The calculation of $\\sigma_e$ that ignored all rotation is, not surprisingly, farthest from the \\citetalias{gul09} line.  Calculation of the relevant quantities for comparison with the $M$-$L$ relation is straightforward, and we plot our value of \\mbhs and $L_V$ (green square) along with that from \\citetalias{gul09} in the right-hand panel.  \\subsection{Globular Clusters} \\label{gcsection}       As demonstrated in section \\ref{orbitstructure}, we find significant radial anisotropy in the globular clusters.  It is interesting that the stellar kinematics at smaller radii do not show this feature.  This difference in orbital properties combined with the difference in their light profiles might suggest the GCs and stars are two distinct populations of tracer particles. This could also indicate the two populations have different formation scenarios.  Unfortunately, there is no radius in the galaxy where we have simultaneous coverage of both stellar and GC kinematics.  Thus, we are unable to test whether the stellar orbits become more radial in the $\\sim 50\\arcsec$ between where the stellar kinematics run out and the GCs begin.  However, since the light profiles of both populations are significantly different (Figure \\ref{allcomps}) there is no reason to assume they should share similar orbit properties.     Figure \\ref{vr_vt} shows the globular clusters in NGC~4594 are radially anisotropic ($\\sigma_r/\\sigma_t >1$) over roughly the radial range $100$-$1000\\arcsec$ (approximately 1-10 $R_e$). Previous studies of the GC systems of galaxies have found their velocity ellipsoids to be isotropic \\citep{cot01,cot03}.  However, these studies used spherical Jeans modeling instead of the more general axisymmetric Schwarzschild code we use.   \\citet{rho04} determine with high confidence that the color distribution of the GC system in NGC~4594 is bimodal.  This may indicate different subpopulations of GCs with different orbital properties that formed at different epochs in the galaxy's history. In our analysis, we make no distinction between red and blue subpopulations.  We use the light profile and kinematics of all available GCs, regardless of color.  However, since we use different sources for our kinematics and photometry data, there is the possibility that each source draws from a different GC subpopulation.  \\subsection{Dark Halo}  The parameters $V_c$ and $r_c$ of our model dark halo imply a central dark matter density of $\\rho_c=0.35 \\pm 0.1 \\, M_{\\odot} \\text{ pc}^{-3}$ Using an improved Jeans modeling technique, \\citet{tem06} model NGC~4594 and find a dark matter halo with central density $\\rho_c=0.033 \\, M_{\\odot} \\text{ pc}^{-3}$, ten times lower than our value. They, however, measure a larger stellar \\fml$_V=7.1 \\pm 1.4$ in the bulge.    We plot the fraction of enclosed mass that is dark matter as a function of half-light radius $R_e$ in Figure \\ref{fracdm}.  At $1$ $R_e$ there is already a roughly $50$-$50$ mix of stars and DM.  Inside of $R_e$ the dark matter still contributes a non-negligible fraction to the total mass content.  In a study measuring dark matter properties in $1.7 \\times 10^5$ local ($z<0.33$) early-type galaxies from the Sloan Digital Sky Survey, \\citet{gri10} find a correlation between the fraction of dark matter within $R_e$ and the logarithmic value of $R_e$. With our measured value of $R_e$, this correlation predicts a dark matter fraction at $R_e$ of $0.68$.  Our value of $0.52$ is smaller, but still within their 68\\% confidence limit.  \\citet{tho09} derive scaling relations for halo parameters based on observations of early-type galaxies in the Coma cluster.  These relations are constructed for similar galaxies using the same halo parameterization and modeling code used in this paper.  This makes comparison to our parameters straightforward. We compare to the observed relations between halo parameters $r_c$, $V_c$, and $\\rho_c$ and total blue luminosity $L_B$.  Our value of $V_c$ falls directly on the $V_c$-$L_B$ relation, however our measured $r_c$ is smaller by roughly an order of magnitude.  Since $\\rho_c \\propto V^2_c / r^2_c$, the discrepancy in $r_c$ causes our measurement of $\\rho_c$ to be high when compared to the \\citet{tho09} $\\rho_c$-$L_B$ relation.  Scatter in this relation is large, however, and the environment of NGC~4594 is different from that of the Coma galaxies.    \\citet{kor04b, kor11b} also derive scaling laws for similar parameters in galaxies of later Hubble type (Sc-Im).  We measure a much higher density and much smaller core radius than the \\citet{kor04b,kor11b} relations imply at the $L_B$ of NGC~4594.  We interpret this as the result of severe compression of the halo by the gravity of the baryons \\citep{blu86}. Such an effect is expected in early-type galaxies with massive bulges."
    },
    "1107/1107.5307_arXiv.txt": {
        "abstract": "We report the discovery by the \\Swift\\ hard X-ray monitor of the transient source \\Swift\\ J2058.4+0516 (\\event).  Our multi-wavelength follow-up campaign uncovered a long-lived (duration $\\gtrsim$ months), luminous X-ray ($L_{\\mathrm{X,iso}} \\approx 3 \\times 10^{47}$\\,erg\\,s$^{-1}$) and radio ($\\nu L_{\\nu,\\mathrm{iso}} \\approx 10^{42}$\\,erg\\,s$^{-1}$) counterpart.  The associated optical emission, however, from which we measure a redshift of 1.1853, is relatively faint, and this is not due to a large amount of dust extinction in the host galaxy.  Based on numerous similarities with the recently discovered GRB\\,110328A / \\Swift\\ J164449.3+573451 (\\eventb), we suggest that \\event\\ may be the second member of a new class of relativistic outbursts resulting from the tidal disruption of a star by a supermassive black hole.  If so, the relative rarity of these sources (compared with the expected rate of tidal disruptions) implies that either these outflows are extremely narrowly collimated ($\\theta < 1^{\\circ}$), or only a small fraction of tidal disruptions generate relativistic ejecta.  Analogous to the case of long-duration gamma-ray bursts and core-collapse supernovae, we speculate that rapid spin of the black hole may be a necessary condition to generate the relativistic component.  Alternatively, if powered by gas accretion (i.e., an active galactic nucleus [AGN]), \\event\\ would seem to represent a new mode of variability in these sources, as the observed properties appear largely inconsistent with known classes of AGNs capable of generating relativistic jets (blazars, narrow-line Seyfert 1 galaxies). ",
        "introduction": "\\label{sec:intro} The recent discovery of the transient source GRB\\,110328A / \\Swift\\ J164449.3+573451 (\\eventb) has unveiled an entirely new class of high-energy outbursts \\citep{ltc+11,bkg+11,zbs+11}.  Like long-duration gamma-ray bursts (GRBs), the outburst was believed to mark the birth of a relativistic jet, generating luminous X-ray and radio emission. However, the central engine powering \\eventb\\ was the super-massive ($M_{\\mathrm{BH}} \\lesssim 10^{7}$\\, M$_{\\odot}$) black hole in the nucleus of an otherwise normal (i.e., nonactive) galaxy \\citep{bgm+11}.  While not conclusive, the observed emission may result from the tidal disruption \\citep{r88} of a star (or white dwarf; \\citealt{kp11}) passing too close to the central black hole \\citep{bgm+11,bkg+11,zbs+11,ctl11}.  A number of tidal disruption flare (TDF) candidates have been previously identified at longer wavelengths \\citep[e.g.,][]{rgd+95,bkd96,kg99,gsz+00,dbe+02, gmm+06,gbm+08,ghc+09,esf+07,esk+08,mue10,vfg+10,cbk+11}.  But the robust inference of a newly born relativistic jet, and the insights into the accretion and jet-formation processes provided therein, clearly distinguish \\eventb\\ from previous TDF candidates.  In this work, we report on a recently discovered high-energy transient, \\Swift\\ J2058.4+0516 (\\event).  Remarkably, \\event\\ shares many of the properties that made \\eventb\\ such an exceptional event, in particular (1) a long-lived (duration $\\gtrsim$ months), super-Eddington ($L_{\\mathrm{X,iso}} \\approx 3 \\times 10^{47}$\\,erg\\,s$^{-1}$) X-ray outburst; (2) a luminous radio counterpart, indicating the presence of relativistic ($\\Gamma \\gtrsim 2$) ejecta; and (3) relatively faint ($M_{U} \\approx -22.7$\\,mag) optical emission.  If future observations are able to establish an astrometric association with the nucleus of the redshift $z = 1.1853$ host galaxy, and also continue to suggest that this galaxy does not harbor an active galactic nucleus (AGN), we may have identified a second member of this new class of relativistic TDFs. Alternatively, several classes of AGNs are known to generate relativistic jets (blazars, narrow-line Seyfert 1 [NLS1] galaxies). If \\event\\ is shown to result from an AGN (e.g., via repeated outbursts over the coming years), we would have uncovered yet another unique example of the already diverse phenomenology of variability from active galaxies.  Throughout this work, we adopt a standard $\\Lambda$CDM cosmology with $H_{0}$ = 71\\,km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\mathrm{m}} = 0.27$, and $\\Omega_{\\Lambda} = 1 - \\Omega_{\\mathrm{m}} = 0.73$ \\citep{sbd+07}. All quoted uncertainties are 1$\\sigma$ (68\\%) confidence intervals unless otherwise noted, and UT times are used throughout. Reported magnitudes are in the AB system \\citep{og83}.  \\begin{figure*}[th] \\plotone{f1.eps} \\caption{Hard X-ray (15--50\\,keV), X-ray (0.3--10\\,keV), and optical light curve of \\event.  The inset in the X-ray panels shows the derived power-law spectral index ($\\Gamma$) as a function of the X-ray count rate (i.e., flux).  Inverted triangles represent $3\\sigma$ upper limits.} \\label{fig:lcurve} \\end{figure*} ",
        "conclusions": "\\label{sec:disc}  \\subsection{Comparison to \\eventb\\ and Known Classes of Active Galaxies} \\label{sec:swj1644} In Figure~\\ref{fig:lxlo}, we plot the X-ray and optical luminosity of \\event\\ at an (observer-frame) time of $\\Delta t \\approx 12$\\,d after discovery.  For comparison, we also plot analogous measurements for a sample of long-duration GRB X-ray afterglows (at the same epoch post-burst), as well as nearby AGNs, more distant, luminous quasars from SDSS, and some of the most dramatically variable blazars \\textit{observed while in outburst}.  With an isotropic X-ray luminosity $L_{\\mathrm{X,iso}} \\approx 3 \\times 10^{47}$\\,erg\\,s$^{-1}$, \\event\\ is much too luminous at this late time to result from a GRB.  Yet with an absolute magnitude of $M_{U} \\approx -22.7$\\,mag, the observed optical emission is several orders of magnitude underluminous compared with the brightest X-ray quasars and blazars.  Even compared with simultaneous X-ray and optical observations of the strongest high-energy flares from some of the best-known blazars (e.g., Markarian\\,501, \\citealt{pvt+98}; 3C\\,279, \\citealt{wpu+98}; Markarian\\,421, \\citealt{fbb+08}), the ratio of X-ray to optical luminosity appears to be quite unique. However, \\event\\ occupies a nearly identical region of this phase space as \\eventb.  \\begin{figure*} \\plotone{f4.eps} \\caption{X-ray and optical luminosity of \\event, compared with a sample of nearby active galaxies (red circles; \\citealt{h09}), luminous quasars from SDSS (cyan circles), galaxies from SDSS with XMM counterparts (black dots; \\citealt{pmc+11}), long-duration GRB afterglows (extrapolated to a common epoch of $\\Delta t = 12$\\,d after the GRB trigger; \\citealt{kkz+07,ebp+09}), and high-energy outbursts from known blazars (Markarian\\,501, \\citealt{pvt+98}; PKS\\,2155-304, \\citealt{fgt+07}; PKS\\,0537-441, \\citealt{prt+07}; 3C\\,279, \\citealt{wpu+98}; Markarian\\,421, \\citealt{fbb+08}). Compared with these sources (and at the same epoch post-discovery as the GRBs), both \\event\\ and \\eventb\\ exhibit extremely luminous X-ray emission, yet relatively faint optical emission.  While for \\eventb\\ this is to some extent due to dust extinction in the host galaxy (for which we have not corrected here), the blue UV/optical SED of \\event\\ suggests at most a modest host-galaxy column density.  We note that we have not K-corrected the absolute magnitudes reported here (aside from cosmological stretch) due to the relatively uncertain intrinsic SEDs.  Adapted from \\citet{ltc+11}.} \\label{fig:lxlo} \\end{figure*}  This similarity is further reinforced when adding the radio observations to the broadband spectral energy distribution (SED).  In Figure~\\ref{fig:sed} we plot the SED of \\event\\ at $\\Delta t \\approx 43$\\,d after discovery, compared with analogous measurements for a sample of blazars, the most dramatically variable class of AGNs (and also known sources of relativistic jets; \\citealt{up95}; see \\S\\ref{sec:physics}).  Blazars exhibit a well-defined luminosity ``sequence'' whereby the lower-frequency maximum of their double-peaked SED occurs at lower energies for more luminous sources \\citep{fmc+98}.  While the radio and optical emission from \\event\\ are comparable to those seen in low-luminosity blazars, the X-ray emission outshines even the brightest sources of this class.  On the other hand, the SED of \\eventb, at a time of $\\Delta t \\approx 20$\\,d, provides a good match to the observed properties of \\event\\ (particularly after correcting for the large but uncertain dust extinction from \\eventb).  While the sequence plotted in Figure~\\ref{fig:sed} may represent the ensemble properties of blazars as a class, individual sources may exhibit quite different behavior.  During an outburst, blazar SEDs can undergo radical changes --- for example, as a result of a rapid brightening in 1997, the low-energy (synchrotron) spectral peak in Markarian\\,501 shifted frequencies by approximately two orders of magnitude \\citep{pvt+98}.  We therefore also compare the observed SED of \\event\\ with that of several high-energy outbursts from known blazars in Figure~\\ref{fig:xrayopt}.  Even when compared with these extreme cases, the ratio of the X-ray to optical flux observed in \\event\\ stands out, being similar only to the broadband properties of \\eventb.  The sole remaining defining characteristic of \\eventb\\ is its association with the nucleus of a nonactive galaxy \\citep{ltc+11,zbs+11}.  At the current time, the observed optical emission from \\event\\ appears to be dominated by the transient, and we have no constraints on the location of the host nucleus from pre-outburst observations.  Future high-resolution observations, when the transient light has faded, should be able to constrain the transient-nucleus offset.  That said, the detection of strong absorption features indicates that the line of sight penetrates a dense region of the host galaxy (as in the case of some GRBs).  Finally, we utilize the optical spectrum to search for evidence of past AGN activity.  Like blazars, NLS1 galaxies are another subclass of AGNs that can sometimes power relativistic jets (e.g., \\citealt{aaa+09g,aaa+09h,fgk+11}; \\S\\ref{sec:physics}).  Similarly, some NLS1 galaxies have been associated with luminous ($L_{\\rm X} \\approx 10^{44}$\\,erg\\,s$^{-1}$) X-ray outbursts (e.g., \\citealt{gbm+95,bpf95}).  In the case of \\eventb, all of the detected emission lines were unresolved, and standard diagnostic diagrams revealed no evidence for ionization from a hidden (i.e., Compton thick) AGN \\citep{ltc+11,zbs+11}.  The larger redshift precludes such an analysis for \\event, as the standard diagnostic lines do not fall in the optical bandpass.  However, we have compiled all of the available spectra from SDSS of known NLS1 galaxies (from the catalog of \\citealt{vv10}) with $z > 0.5$ (478 objects), and used these to create an NLS1 template.  This composite spectrum is plotted alongside that of \\event\\ in Figure~\\ref{fig:spec}. The template exhibits strong, broad emission from \\ion{C}{2}] $\\lambda$2324, [\\ion{Ne}{4}] $\\lambda \\lambda$2422, 2424, and in particular \\ion{Mg}{2} $\\lambda \\lambda$2796, 2803.  None of these emission lines are detected in our spectrum of \\event.  In fact, of the 478 known NLS1 galaxies in SDSS with $z > 0.5$, all but one have well-detected ($> 3\\sigma$) \\ion{Mg}{2} in emission.  Unless exceedingly rare, the optical spectrum of \\event\\ suggests that it is not an NLS1 galaxy.  We have furthermore performed a similar analysis with all known BL Lac sources from \\citet{vv10} at $z > 0.5$ in the SDSS archive (68 sources); the resulting median spectrum is shown in red in Figure~\\ref{fig:spec}. By definition, BL Lacs lack the bright emission features present in the spectra of NLS1 galaxies, and therefore they more closely resemble the observed spectrum of \\event.  However, BL Lac spectra do typically exhibit \\ion{Ca}{2} $\\lambda \\lambda$3934, 3968 in absorption, which is lacking in \\event, and not a single BL Lac in the SDSS database exhibits strong \\ion{Mg}{2} absorption.  While we cannot entirely rule out a new mode of variability in a blazar, the unique SED and optical spectrum distinguish \\event\\ from known members of this class as well.  \\begin{figure*} \\plotone{f5.eps} \\caption{SEDs of \\event\\ and \\eventb, compared with the blazar ``sequence'' from \\citet{fmc+98}.  The open squares attempt to correct for the large (but uncertain) host-galaxy extinction in the case of \\eventb\\ by assuming $A_{V} = 5$\\,mag.  The large X-ray luminosities of both \\event\\ and \\eventb\\ are incompatible with the observed optical and radio fluxes, which would imply an intrinsically fainter source. Both objects have complex SEDs that are difficult to reconcile with only a single emission component (e.g., synchrotron).} \\label{fig:sed} \\end{figure*}  \\subsection{Basic Physical Properties} \\label{sec:physics} We have shown in the previous section that \\event\\ does not appear consistent with any known class of AGNs (particularly blazars and NLS1).  Based on the many similarities with \\eventb,  we can use the same arguments to derive the basic physical parameters for \\event\\ (e.g., \\citealt{bgm+11}). From our NIR observations, we derive an upper limit on the host-galaxy absolute $V$-band magnitude of $M_{V} \\gtrsim -21$\\,mag (we have used an S0 galaxy template from \\citealt{kcb+96} for the K-correction based on the lack of emission features in the optical spectrum).  Applying the black hole mass vs. bulge luminosity \\citep{mtr98} correlation from \\citet{lfr+07} and assuming a bulge-dominated system, we limit the mass of the putative central black hole to be $M_{\\mathrm{BH}} \\lesssim 2 \\times 10^{8}$\\, M$_{\\odot}$.  Separately, the shortest observed X-ray variability time scale can be used to constrain the central black hole mass using causality arguments (e.g., \\citealt{GCN.11851}).  For an observer-frame variability time scale of $\\delta t \\lesssim 10^{4}$\\,s, we find $M_{\\mathrm{BH}} \\lesssim 5 \\times 10^{8}$\\, M$_{\\odot}$.  The most dramatic short time scale variability observed from \\eventb\\ occurred at $\\Delta t \\lesssim 10$\\,d \\citep{bgm+11}, before we commenced X-ray observations of \\event.  But even at late times, the X-ray light curve of \\event\\ exhibits a degree of variability (changes of an order of magnitude on time scales of a few days) not seen in \\event, suggesting an imperfect analogy between these two sources.  Finally, we can constrain the central mass using the black hole fundamental plane, a relationship between mass, X-ray luminosity, and radio luminosity (e.g., \\citealt{gcm+09}).  Application of this relation requires knowledge of the beaming fraction, for which we adopt a value of $f_{b} \\equiv (1 - \\cos \\theta) \\approx 10^{-2}$ (see below). Following \\citet{mg11}, we find $M_{\\mathrm{BH}} \\approx 5 \\times 10^{7}$\\, M$_{\\odot}$; however, we caution that the scatter in this relationship is even larger than that found in the black hole mass vs. bulge luminosity relation.  The derived black hole mass, $M_{\\mathrm{BH}} \\lesssim 10^{8}$\\, M$_{\\odot}$, is of vital importance for two main reasons. First, the tidal disruption of a solar-type star will only occur outside the event horizon for $M_{\\mathrm{BH}} \\lesssim 2 \\times 10^{8}$\\, M$_{\\odot}$.  As such, if the origin of the accreting material is indeed the tidal disruption of a nondegenerate star (or more massive object), the mass of the central black hole in the host galaxy of \\event\\ must be sufficiently small\\footnote{The limits on the central black hole of the host galaxy of \\eventb\\ are even more stringent ($M \\lesssim 10^{7}$\\, M$_{\\odot}$; \\citealt{ltc+11,bkg+11}), and may even allow for the disruption of a white dwarf \\citep{kp11}.}.  Future observations of the host, when the optical transient light has faded, should provide a significantly improved estimate of $M_{\\mathrm{BH}}$.  Second, for $M_{\\mathrm{BH}} \\lesssim 10^{8}$\\, M$_{\\odot}$, the observed X-ray emission exceeds the Eddington luminosity.  In particular, the peak isotropic 0.3--10\\,keV X-ray luminosity of $L_{\\mathrm{X,iso}} \\approx 3 \\times 10^{47}$\\,erg\\,s$^{-1}$ exceeds the Eddington value by more than an order of magnitude.  Given the long-lived nature of the central engine, this suggests the outflow is likely to be collimated with a beaming factor $f_{b} \\lesssim 10^{-1}$ (cf., \\citealt{s11}).  An even larger degree of collimation is required to reconcile the observed rate of these outbursts with theoretical predictions for TDFs (\\S\\ref{sec:implications}).  More robust evidence for beaming is provided by the large radio luminosity.  It is well known that the brightness temperature of an incoherent radio source cannot exceed $10^{11}$\\,K for extended periods of time \\citep{r94,kfw+98}.  Utilizing this fact, we can derive a lower limit on the angular size of the radio-emitting region, $\\theta \\gtrsim 14$\\,$\\mu$as.  Assuming that the outflow commencement coincides with the hard X-ray discovery, this implies a mean expansion velocity of $\\beta \\equiv v / c \\geq 0.88$, or a mildly relativistic expansion with $\\Gamma \\geq 2.1$.  Furthermore, the radio emission will be collimated into a viewing angle $\\theta \\lesssim 1 / \\Gamma$ due to relativistic beaming effects.  If we integrate over the X-ray light curve to date (2011 July 20), we find that \\event\\ has emitted a total 0.3--10\\,keV fluence of $S_{\\rm X} \\approx 1.0 \\times 10^{-4}$\\,erg\\,cm$^{-2}$.  If we assume this accounts for $\\sim 1/3$ of the broadband luminosity, the total isotropic energy release is $E_{\\mathrm{iso}} \\approx 10^{54}$\\,erg, comparable to that of \\eventb\\ at a similar epoch.  For a typical radiative efficiency for accretion power ($\\eta \\equiv E_{\\mathrm{rad}} / M c^{2} \\lesssim 0.1$), the tidal disruption of a solar-type star would require a significant beaming correction.  If only a small fraction of the available energy is allocated to the observed relativistic component ($\\lesssim 1$\\%), this would imply $f_{b} \\lesssim 10^{-3}$.  In addition to the weaker degree of X-ray variability, we wish to draw attention to two other important differences between these sources. Unlike that case of \\eventb, the observed UV/optical SED of \\event\\ is quite blue.  In fact, the rest-frame UV data for \\event\\ can be well fit by a blackbody spectrum (Fig.~\\ref{fig:xrayopt}), with $T_{\\mathrm{BB}} \\gtrsim 6 \\times 10^{4}$\\,K, $L_{\\mathrm{BB}} \\gtrsim 10^{45}$\\,erg\\,s$^{-1}$, and $R_{\\mathrm{BB}} \\gtrsim 10$\\,AU.\\footnote{Because the observed filters fall on the Rayleigh-Jeans tail, our constraints on the blackbody temperature, luminosity, and radius are essentially lower limits.} The derived blackbody parameters are similar to the thermal emission observed from previous TDF candidates discovered at UV and optical wavelengths \\citep{gmm+06,gbm+08,ghc+09,vfg+10,cbk+11}, and are strongly suggestive of an accretion disk origin for the optical light.  In this sense, \\event\\ may serve as a link between the bright, nonthermal X-rays observed in the relativistic TDF candidates, and the thermal disk emission from ``classical'' TDFs.  Such comparisons were not possible for \\eventb\\ due to dust in the host galaxy.  More importantly, the radio spectral index of \\event\\ appears to be quite flat ($f_{\\nu} \\propto \\nu^{0}$).  This stands in stark contrast to \\eventb, which was optically thick ($f_{\\nu} \\propto \\nu^{1.3}$) up to high frequencies in the days and weeks following discovery \\citep{zbs+11}.  This may simply reflect physical differences in the jet (e.g., microphysics) or its environment (e.g., density), as is suggested by the different extinction properties.  However, it may also imply a more extended (and hence long-lived) radio source, which would be difficult to reconcile with a tidal disruption origin. Future high-resolution radio interferometry (i.e., VLBI) may be able to resolve this issue.  \\begin{figure*} \\plotone{f6.eps} \\caption{Optical and X-ray SEDs of \\event\\ and \\eventb, compared with some of the most dramatically variable blazars \\textit{while in outburst} (Markarian\\,501, \\citealt{pvt+98}; PKS\\,2155-304, \\citealt{fgt+07}; 3C\\,279, \\citealt{wpu+98}).  Unlike any of the other objects plotted here, the rest-frame UV SED of \\event\\ is well represented by a blackbody by $T \\gtrsim 6 \\times 10^{4}$\\,K (solid black line).} \\label{fig:xrayopt} \\end{figure*}"
    },
    "1107/1107.2657_arXiv.txt": {
        "abstract": "% { As a galaxy evolves, its chemical composition changes and the abundance ratios of different elements are powerful probes of the underlying evolutionary processes. Phosphorous is an element whose evolution has remained quite elusive until now, because it is difficult to detect in cool stars. The infrared weak \\ion{P}{i} lines of the multiplet\\,1, at 1050-1082\\,nm, are the most reliable indicators of the presence of phosphorus. The availability of CRIRES at VLT has permitted access to this wavelength range in stellar spectra. } { We attempt to measure the phosphorus abundance of twenty cool stars in the Galactic disk. } { The spectra are analysed with one-dimensional model-atmospheres computed in Local Thermodynamic Equilibrium (LTE). The line formation computations are performed assuming LTE. } {The ratio of phosphorus to iron behaves similarly to sulphur, increasing towards lower metallicity stars. Its ratio with respect to sulphur is roughly constant and slightly larger than solar, [P/S]=$0.10\\pm 0.10$. } {We succeed in taking an important step towards the understanding of the chemical evolution of phosphorus in the Galaxy. However, the observed rise in the P/Fe abundance ratio is steeper than predicted by Galactic chemical evolution model model developed by Kobayashi and collaborators. Phosphorus appears to evolve differently from the light odd-Z elements sodium and aluminium. The constant value of [P/S] with metallicity implies that P production is insensitive to the neutron excess, thus processes other than neutron captures operate. We suggest that proton captures on $^{30}$Si and $\\alpha$ captures on $^{27}$Al are possibilities to investigate. We see no clear distinction between our results for stars with planets and stars without any detected planet. } ",
        "introduction": "Phosphorus is an abundant element in the Universe, and in the solar photosphere it is among the top twenty most abundant elements. In the periodic table of elements, phosphorus, with Z=15, is in the same group as nitrogen, lying between silicon (Z=14) and sulphur (Z=16). Silicon is a well-studied element, and its abundance is easily measured in photospheres  of stars of type F, G, or K, by analysing the \\ion{Si}{i} lines. In the last fifteen years, sulphur has been systematically investigated by several groups that analysed the few multiplets of \\ion{S}{i} available in the observed spectra (for further details see \\citealt{spite11} and references therein). This is not the case for phosphorus, that, before this work, had never been analysed systematically in cool stars. The reason why was already given by \\citet{struve30}: no \\ion{P}{i} line is available in the ``ordinary'' range of the observed spectra of stars of spectral type F, G, or K. Some \\ion{P}{ii} and \\ion{P}{iii} lines are observable in the spectra of B-type stars. In the UV spectrum, some lines of  \\ion{P}{iv}, \\ion{P}{iii}, and \\ion{P}{ii} can be detected in hot stars. To study the variation in the phosphorous abundance over time we need to use long-lived stars of type F, G, or K. In the era of the high resolution spectrograph CRIRES-VLT, some lines of \\ion{P}{i}, the lines of Mult.\\,1, at 1050-1082\\,nm, can be observed. For the first time, it has thus become possible to trace the Galactic evolution of phosphorus.  One single stable isotope of phosphorus, $^{31}$P, is believed to be formed via neutron capture, as for the parent nuclei $^{29}$Si and $^{30}$Si, probably in the carbon and neon burning shells during the late stages of the evolution of massive stars. The $^{31}$P produced in this way is then released by the explosion of these massive stars as type II SNe. \\citet{woosley95} expect no significant production of phosphorus during the explosive phases.  The handful of observations available of P in different stars was summarised in \\citet{caffau07}. Since then, \\citet{hubrig09} have derived the P abundance in horizontal branch (HB) stars in globular clusters NGC 6397 and NGC 6752, where it appears to be strongly enhanced, by more than 2 dex and almost 3 dex, respectively. Chemical anomalies in HB stars are usually attributed to diffusion effects, but the current models by \\citet{michaud08} do not seem to account for more than about 0.7\\,dex of the P enhancement in the photosphere of HB stars of similar temperature.  \\citet{melendez09} measured P in ten solar twins, although they do not state which lines they used. The lines of Mult.\\, 1 were not available in their observations and we speculate that they used the strongest  line of Mult.\\,2 at 979.6\\,nm.  A result concerning P evolution that we find intriguing, is that of \\citet{sbordone09}, who found that the sulphur abundance in the globular cluster 47 Tuc displays a significant correlation with the Na abundance. The authors point out that S could be produced through proton capture reactions on P ($^{31}{\\rm P}\\left( p,\\gamma\\right)^{32}{\\rm S}$). However, P in the Sun is roughly 1.7 dex less abundant than S. If the same abundance ratio held elsewhere, it would be impossible to produce any significant amount of S at the expense of P. The P abundance has never been measured in stars of 47 Tuc.  In this paper, we use high-resolution infrared (IR) spectra, obtained with CRIRES at VLT, to analyse the behaviour of the P abundance as a function of metallicity in the Galactic disc. We selected a sample of stars from \\citet{eczolfo} for which the S abundance is known, to assess whether the proton capture on P is a viable production channel for S. The sample spans a metallicity range of 1.2\\,dex, $-1.0<{\\rm\\left[Fe/H\\right]}<+0.3$. Our future plan is to extend the sample of metal-poor stars, and to include metallicities lower than [Fe/H]\\footnote{ [X/H]$=\\log\\left({\\rm N_X/N_H}\\right)-\\log\\left({\\rm N_X/N_H}\\right)_\\odot$ and [X/Y]=[X/H]-[Y/H]} =--1. ",
        "conclusions": "We have determined the phosphorus abundance in twenty stars of the Galactic disk that we observed with CRIRES at VLT. The metallicity of the sample extends over about 1.2\\,dex, providing a first insight into the evolution of phosphorus in the Galaxy. The only published model of the chemical evolution of the Galaxy that provides information about the evolution of phosphorus is that of \\citet{kobayashi06}. The model predicts correctly, from a qualitative point of view, the rise of the P/Fe ratio with decreasing metallicity. The observed rise is, however, larger by almost a factor of two than the model's prediction, and lacks the ``plateau''' at metallicities [Fe/H] $> -0.5$. Our analysis indicates that phosphorus behaves in a similar way to an $\\alpha$-element, such as sulphur. The model of Galactic chemical evolution described by \\citet{Gibson} predicts a $\\sim 0.2$\\,dex enhancement of the [P/Fe] ratio over two dex in metallicity, from solar to $-2.0$. The level of enhancement at low metallicity is compatible with what we observe, but no mention is made of the decrease in this ratio near solar metallicity.   We have found that the [P/S] is roughly constant at a level of about 0.1\\, dex. However, the star-to-star scatter makes this value compatible with zero. This constant value implies that P and S are produced in the same amounts over the metallicity range considered by our investigation. The picture is different from what is found for Na or Al: both [Na/Mg] and [Al/Mg] decrease significantly with decreasing metallicity. A possible interpretation is that the P production is insensitive to the neutron excess $\\eta$. If this is the case, then P should be produced by nuclear reactions other than neutron captures. Possibilities that should be explored are proton captures on $^{30}$Si and $\\alpha$ captures on $^{27}$Al. A corollary of this analysis is that, since in the Sun sulphur is about a factor of 30 more abundant than phosphorus, this ratio remains the same at all metallicities (in the range that we have examined). We therefore conclude that S production through proton captures on P is not a viable mechanism. Although the possibility of a large enhancement of P abundance in globular clusters cannot be ruled out, it seems unlikely.  It is interesting to compare our results to phosphorus abundances derived from measurements in damped Ly-$\\alpha$ (DLA) systems at high redshift. These are gas-rich galaxies observed as absorption lines in the direction of QSOs. \\citet{Molaro} measured the P abundance in the DLA at $z=3.3901$ towards QSO 0000-2620. This system has a very low metallicity and measurements of [P/H]=$-2.31$, [P/Fe]=$-0.27$, and [P/S]=$-0.40$. The ratios of P to iron and sulphur are clearly different from what those seen in Galactic disk stars. This may be  due to the low metallicity of the system, suggesting that perhaps also in our Galaxy, the [P/Fe] ratio starts at low values, increases until reaching a maximum, and thereafter decreases towards solar metallicity. This behaviour is predicted by the model of \\citet{kobayashi06}, although not shown in Fig\\,.2, which only extends down to [Fe/H]=--1.0. Another possibility is that this DLA galaxy has undergone a different chemical evolution, since it has [S/Fe]=+0.13, while at these metallicities [S/Fe]$\\sim +0.4$ in the Galaxy \\citep{eczolfo}.  \\citet{Fenner} determined an upper limit of [P/S]$<+0.01$ in the DLA at $z=2.626$ towards FJ081240.6+320808. This system is more metal-rich ([O/H]=$-0.44$), thus probably of metallicity comparable to some of the stars in our sample. However, the [P/S] ratio is slightly lower than what we observe. Although this DLA system may resemble the Milky Way disk, it has probably experienced a different chemical evolution.  \\balance  Our investigation is the first step towards an understanding of the evolution of P in the Galaxy. It will be important to secure more data at metallicity [Fe/H]=$-1.0$ and lower, to understand whether the [P/Fe] ratio flattens out, continues to rise, or finally decreases towards low metallicity. The good performance of the CRIRES spectrograph suggests that it will be possible to observe considerably fainter targets with adequate exposure times. A direct measurement of P in 47 Tuc or other globular clusters would allow us to probe the feasibility of S production through proton capture in these objects."
    },
    "1107/1107.0187_arXiv.txt": {
        "abstract": "Large-scale bars and minor mergers are important drivers for the secular evolution of galaxies. Based on ground-based optical images and spectra as well as ultraviolet data from the {\\it Galaxy Evolution Explorer} and infrared data from the {\\it Spitzer Space Telescope}, we present a multi-wavelength study of star formation properties in the barred galaxy NGC 7479, which also has obvious features of a minor merger. Using various tracers of star formation, we find that under the effects of both a stellar bar and a minor merger, star formation activity mainly takes place along the galactic bar and arms, while the star formation rate changes from the bar to the disk. With the help of spectral synthesis, we find that strong star formation took place in the bar region about 100 Myr ago, and the stellar bar might have been $\\sim$10 Gyr old. By comparing our results with the secular evolutionary scenario from Jogee et al., we suggest that NGC 7479 is possibly in a transitional stage of secular evolution at present, and it may eventually become an earlier type galaxy or a luminous infrared galaxy. We also note that the probable minor merger event happened recently in NGC 7479, and we find two candidates for minor merger remnants. ",
        "introduction": "The physical processes of galaxy evolution can be classified into fast and slow according to their timescales, and also can be divided into internal and external based on origins of drivers \\citep{KK04}. As the Universe expands, internal secular processes will become important. Secular evolution is different from dissipative collapses and mergers of galaxies, which are rapid and violent \\citep[e.g.,][]{Sandage90,Sandage05}. Secular processes can make the energy and mass rearrange, and be driven by galactic structures such as bars, oval disks, spiral arms, triaxial dark halos inside of galaxies \\citep{Athanassoula02, KK04, Kormendy05, Kormendy08}, also by minor mergers and other environmental factors outside the galaxies \\citep{Bournaud05, Bournaud07, Jogee06}.  Many observations and researches have shown that at low redshift, especially in the local universe (z $\\sim$ 0), mergers are less common \\citep[e.g.,][]{Le00, ConseliceB03}, so secular evolution will have important effects on galaxies \\citep{KK04}. Pseudobulges, galactic structures produced by secular evolution, are common in disk galaxies \\citep{Kormendy05, Kormendy08}, and also indicates that this evolution is a common phenomenon to galaxies. This process is not only limited to low redshift galaxies, e.g., \\citet{Genzel08} found that secular evolution may help account for the stellar buildup observed in massive galaxies at z $\\sim$ 2. In addition, as the Universe expands, secular evolution will take the place of mergers to dominate the evolution of galaxies \\citep{KK04}.  As important internal structures to secular evolution, galactic stellar bars have been discussed in many works, both theoretical and numerical \\citep[e.g.][]{Sellwood99, Barazza07, Athanassoula09}. The bar is a common phenomenon in the universe because about 65$\\%$ of nearby spiral galaxies have bars, among which over 30$\\%$ have strong bars \\citep{Eskridge00}, and the percentage of bars remains high out to redshift z $\\sim$ 1 \\citep{Elmegreen04, Jogee04}. Bars can provide intense nonaxisymmetry in the gravitational potential, which will make gas lose angular momentum and then fall down to the inner region of galaxies \\citep{Athanassoula03, Kormendy05}. \\citet{Sheth05} compared the molecular gas distribution in a sample of nearby galaxies from the BIMA CO (J=1-0) Survey of Nearby Galaxies (SONG), and found that more molecular gas is concentrated in the central kpc of barred spirals than that in other Hubble type galaxies, consistent with radial inflow driven by the bar. After the gas density becomes supercritical in the central parts, the next step is to trigger circumnuclear star formation \\citep{Sheth05}, and (pseudo)bulges would also grow in this progress \\citep{KK04}. These have been proved by a large number of observations \\citep[e.g.,][]{Knapen06, Shi06, Gadotti10} and simulations \\citep[e.g.,][] {Athanassoula09}. \\citet{Gadotti01} also studied the broadband UBV color profiles for 257 Sbc barred and nonbarred galaxies, and found the bulges of barred galaxies are bluer than others, indicating an increase of the star formation rate (SFR) in the central regions of these objects. \\citet{Regan06} found four barred galaxies of 11 spiral galaxies have strong central excesses in both 8 $\\mu$m and CO emission. Whereas \\citet{Ho97a} found that the same effect only in early-type barred spirals using the luminosity and the equivalent width of H$\\alpha$ emission. \\citet{Fisher06} found that some barred galaxies don't have higher SFRs than those of other types in a sample of 50 galaxies spanning Hubble types E to Sc using Spitzer infrared color profiles.  Similar to stellar bars, minor mergers also can drive gas into the galactic central region and fuel nuclear activities \\citep{Jogee06}. Observations and N-body simulations have shown that minor mergers can lead to high SFR \\citep{Ferreiro04, Ferreiro08, Kaviraj09}, engender circumnuclear rings \\citep{Knapen04, Mazzuca06}. In addition, the evolution of stellar bars may relate to minor mergers in numerical simulations \\citep{Berentzen03, Emilio08}. However, these correlations still need to be studied further.  In order to better understand the effects of bars on the secular evolution of galaxies, it is essential to study the star formation properties of nearby barred galaxies. Thus we selected a barred galaxy sample from three large Spitzer Legacy or GTO programs: the Spitzer Infrared Nearby Galaxies Survey \\citep[SINGS;][]{Kennicutt03}, the Local Volume Legacy Survey \\citep[LVLS;][]{Lee08}, and the Infrared Hubble Atlas \\citep{Pahre04}. It is consisted of 73 nearby (z $<$ 0.02) barred galaxies (including Hubble types SB \\& SAB). One most interesting object in this sample is NGC~7479. It is a SBbc galaxy \\citep{de91} with a large stellar bar and two strongly asymmetric spiral arms, among which the western one is much stronger. Its small and bright nucleus is classified as LINER \\citep{Keel83} and Seyfert 1.9 \\citep{Ho97b}. Observations and simulations have been made in the studies of NGC~7479 \\citep[e.g.,][]{Martin97, LaineH99, Aguerri00, Huttemeister00, Laine06, Laine08}. There are amount of molecular gas along the stellar bar and in the central region of NGC~7479, and also obvious dust lanes and large-scale shocks along the bar \\citep{Laine98}. \\citet{Rozas99} showed that intense star formation regions can also be found in the barred region. \\citet{Huttemeister00} found that there may be a nuclear ring, a torus, disk or inner bar in the bulge of this galaxy. It seems that this galaxy has no close companions although it looks like an interacting system from its morphology \\citep{Laine98, Saraiva03}, so it may have suffered a minor merger recently \\citep{Laine08}. Based on the facts above, the secular evolution driven by bar and minor merger simultaneously make NGC~7479 an interesting galaxy worth careful study.  In the present paper, we study the properties of NGC~7479, mainly focusing on the bar and star formation regions based on multi-wavelength data from ultraviolet (UV) to infrared (IR). It is organized as follows: Section 2 describes the data we used and their reductions. Section 3 presents the main results from this study. Section 4 gives some discussions, and we made a summary in section 5. ",
        "conclusions": "NGC 7479 is an interesting example with the strong stellar bar, and have likely suffered a minor merger because of some features mentioned above. The active star formation and distinctive morphology are evidences of the secular evolution driven by the stellar bar and maybe also by the minor merger event. In this section, we will discuss the effects of both drivers to our galaxy, respectively.  \\subsection{Bar Driven Star Formation and Galactic Evolution} \\label{discussion_bar} Based on what we have obtained above, the SFR in the bulge region of NGC7479 is $\\sim$ 1 $M_{\\odot}yr^{-1}$. The star formation activity in this region is likely related to the presence of the large-scale bar. Theoretically, stellar bars can concentrate a substantial fraction of disk gas in small galactocentric radii as a result of star formation \\citep{KK04}. Observations also find that galaxies with stellar bars show higher central SFRs than unbarred ones on average \\citep[e.g.,][]{Kormendy05, Fisher06, Shi06}, and the star formation in circumnuclear regions of spiral galaxies are more dependent on stellar bars than that in their outer disks \\citep{Kennicutt98}. Using mid-IR color to trace star formation activity, \\citet{Roussel01} also found that the color distributions of strong barred galaxies show 15 $\\mu$m excesses, which are absent from weakly barred or unbarred galaxies. Using the specific SFR in the galactic central region of NGC 7479 (Figure~\\ref{fig6}), we estimated that the growth time of the bulge at present SFR is about 50 $\\sim$ 100 Myr, while \\citet{Fisher09} derived that the median pseudobulge could have grown the present-day stellar mass in 8 Gyr, which is much longer than that of NGC 7479. Therefore, there is probably another factor affecting the evolution of this galaxy.  Besides the bulge, star formation activity is also located in the galactic bar of NGC 7479. \\citet{Martin97} studied the morphological properties of H {\\sc ii} regions in the bar of NGC 7479 using its H$\\alpha$ image and found that the total SFR from H {\\sc ii} regions in the bar excluding nucleus and circumnuclear region is about 0.4M$_{\\odot} $yr$^{-1}$. \\citet{LaineK99} also found the SFR is about 0.5M$_{\\odot} $yr$^{-1}$ in the bar region, though this value is lower than what we have estimated using composite tracers. Previous studies also revealed that there is a large mount of gas distributing along the bar. Using CO (J=1 $\\to $ 0) observations, \\citet{LaineK99} found that CO emission in NGC 7479 is along the dust lane traversing the whole bar and nucleus. \\citet{Huttemeister00} also found that a continuous gas distribution is all along the bar in NGC 7479 and the distributions of ${}^{12}CO$, ${}^{13}CO$, dust lanes and velocity gradient have complex relationships. The radio continuum emission at 21 cm in the galactic central region shows the similar behavior as the CO emission \\citep{Laine98}. The distribution of star formation activity and gas in the bar may be a result of the concentration of gas inflowing from the disk under the influence of the bar, because the gravitational torque of the large-scale bar can make the gas lose angular momentum, and then drive it down inward through the stellar bar \\citep{Athanassoula92, Sellwood93}.  In general, star formation locates in the circumnuclear regions and in the ends of bars, while bars are typically dominated by evolved stellar populations \\citep{Gadotti06}. In the study of the barred spiral NGC 5383, \\citet{Sheth00} found weak H$\\alpha$ emission in the bar between the bar ends and the central region, despite the high gas column density in the bar dust lanes. Similarly, weak star formation in the bar takes place in NGC 1530, probably due to the inhibition of the shear and shock in the dust lanes \\citep{Reynaud98}. However, as Figure~\\ref{fig1} shows, obvious H$\\alpha$ emission takes place in the bar of NGC 7479, and a majority of young stars are formed in nearly 100Myr (see Table~\\ref{tabl3}). In addition, the SFR in the bar region is also comparable with those in the bulge and bar ends (Figure~\\ref{fig6}). For that reason, the existence of star formation activity appears to be consistent with not only the galactic bar, but also another effect (see the next subsection).  How does the bar affect the evolution of NGC 7479? \\citet{Quillen95} found that the gas inflow rate along the stellar bar is about 4--6 M$_{\\odot}$yr$^{-1}$ in this galaxy, which is much larger than the SFR in the bar region we derived (1.24 M$_{\\odot}$yr$^{-1}$). Therefore, before turning into stars a large percent of gas will drop into the galactic central region and make the bulge mass addition \\citep{LaineK99}. In this case, this galaxy will possibly evolve toward an earlier galaxy type, although there is a doubt that a late-type galaxy maybe not evolve into an early type via bar-induced gas inflow \\citep{Sheth05}. Therefore, NGC 7479 is likely just in a transitional stage.  In order to explore the bar-driven secular evolution, \\citet{Jogee05} classified 10 different barred galaxies into three evolutionary stages based on the dynamical properties and dust morphology: (i) Type I Non Starburst, the early stages of bar-driven inflow, where large amounts of gas are along the bar and there is low star formation efficiency in the inner few kiloparsecs; (ii) Type II Non Starburst, the later stages of bar-driven inflow where most molecular gas is concentrated in inner kpc and star formation only occurs in regions with high gas densities; subsequently, (iii) Circurmnuclear Starburst, where large fraction of molecular gas is concentrated in galactic central region, exists intense star formation activity. Comparing our results with this scenario suggests that NGC 7479 may be in Type I Non Starburst evolutionary phase based on the properties of star formation and molecular gas in this galaxy. NGC 7479 is similar to NGC 4569, which is the prototype of Type I Non Starburst in the work of \\citet{Jogee05}. With a sample of barred galaxies, \\citet{Verley07} also put forward a scenarios of bar-driven secular evolutionary sequence using properties of H$\\alpha$ emission, in which they defined four main groups, i.e., Group E, a strong central peak in the H$\\alpha$ emission, no H$\\alpha$ emission in bar; Group G, H$\\alpha$ emission in the bar; Group EG, a transition between the E and G groups; Group F, with less gas, smoother morphology and no central emission spot in H$\\alpha$, and argued that the evolutionary sequence is G $\\to $ EG $\\to $ E $\\to $ F. In their studies, NGC 7479 is classified as Group G. Similar results are also found by \\citet{Martin97}, our galaxy suggested to be in the middle transitional type in their three stages of evolutionary sequence.  Thus, as performed by above authors, NGC 7479 is in a temporal transitional phase. Under the effect of its stellar bar, most of molecular gas will be concentrated in the inner kpc of the galactic center, starburst can be triggered in circumnuclear region when most of the gas exceeds a critical density, and pseudobulges is expected to be built in the inner kpc. Then NGC 7479 may become a galaxy with a earlier Hubble type. On the other hand, the increase of the infrared luminosity may also make our galaxy a luminous infrared galaxy (LIRG), this may explain why there are some LIRGs are isolated barred galaxies but not interacting systems \\citep[e.g.,][]{Wang06}.  \\subsection{Minor Merger Event in NGC7479} Minor merger event is a common phenomenon in the Universe. According to statistics, each field spiral galaxy has more than one companions on average \\citep{Zaritsky97}. Thus, the interactions between galaxies and their companions are believed to be very common \\citep{Ostriker75}. Several properties of NGC 7479 indicate the presence of a recent minor merger. First of all, as we mentioned in Sec~\\ref{discussion_bar}, star formation activity in the bar and bulge of our galaxy is stronger than most common barred galaxies. The systems in minor mergers have higher SFR than those found in normal H {\\sc ii} regions of spiral galaxies \\citep{Ferreiro08}. Minor mergers also likely play a role in enhancing the infrared luminosity for some low-luminosity galaxies \\citep{Kartaltepe10}. These features are likely caused by a large amount of gas which was driven into host galactic disks when they swallowed gas-rich satellites, and the gas may also be transported into the galactic inner regions due to bars.  Second, in numerical simulations, vertical impact of small companions can cause asymmetries to gas-rich barred galaxies \\citep{Berentzen03}. \\citet{Walker96} found that minor merger can form significant disk asymmetry, visible for at least 1 Gyr in their simulations. Similarly, the simulations of \\citet{Bournaud05} indicate that minor mergers in some conditions can create a strong asymmetry when the galaxy appears to be ``isolated'', once the most obvious merger related features have disappeared. Through comparing the distribution of different compositions in Figure~\\ref{fig1}, we found that star formation regions of NGC 7479 are largely concentrated in the strong arm, and PAH emission is mostly distributed along the stellar bar and the strong spiral arm. Probably because of this reason, the asymmetry of NGC 7479 is much higher in UV, mid-infrared, as well as H$\\alpha$ band (A $\\sim$ 0.3), than that in the optical and near-infrared bands.  Third, minor mergers are probably related to the triggering of stellar bars \\citep{Shumakova05, Emilio08}. Since active star formation exists in the bar of NGC 7479, the intense gas concentration and young stars in the stellar bar are probably driven by the minor merger event in some degree besides the stellar bar. The large percentage of burst population in the bar of NGC 7479 (Table~\\ref{tabl3}) are evidently formed in recent hundred Myr. Furthermore, as argued by \\citet{Martin00} there is strong evidence that the present morphology of NGC 7479 was acquired following a merging event with a small galaxy about 300 Myr ago. In addition, the radio continuum jet in NGC 7479 found by \\citet{Laine08} is another signal of recent perturbation by a minor merger in this galaxy, and is similar with the cause of jets in NGC 1097 \\citep{Higdon03}. Through the comparison between the simulation model of minor merger and observation, \\citet{LaineH99} suggested that NGC 7479 is the result of minor merger which is still under process, and the remnant of the satellite galaxy likely locates within the bar. Besides the minor merger, there are likely other processes to contribute large amounts of gas to the evolution of NGC 7479. For example, galaxy tidal encounters \\citep{Li08} and external gas accretion \\citep{Bournaud02} that can also result in material infalling; in the simulation of \\citet{Sellwood99}, a tidal encounter or the accreted low angular momentum material could also strengthen the bar.  Although morphological properties of NGC 7479 are consistent with characteristics of minor merger, no companions have been found in the field of this galaxy \\citep{Laine98, Saraiva03}. To explore whether any remnant of merging companions is still existing, we compared 0.2--12.0 keV X-ray images from XMM-Newton telescope, high-resolution images from HST Proposal 6266 with WFPC2 and the F814W filter. Since ultraluminous compact X-ray sources are probably dwarf companion galaxy nuclei \\citep{Makishima00, Wu02}, and merger remnants if still exist are also likely located in compact optical sources, we try to find possible candidates from these sources.  In Figure~\\ref{fig9} we marked 10 candidates, among which there are 4 strong X-ray emission peaks (labeled as X1-X4) detected in X-ray image and 5 intensive optical regions (labeled as O1-O5) in HST image, the last one is a potential target pointed by \\citet{LaineH99} (marked with L). Besides those, the nucleus of the host galaxy (marked as C) can also be clearly detected by XMM images. Target X4 is the nearest ($\\sim$1kpc) X-ray emission peak to the galactic nucleus C, and it lies in the most obscured region, which may be the reason it can't be identified in optical wavelength, and distinguished from the nucleus. Therefore, we could not do the further investigations for this source. Target X1 located in the north end of the bar, has the most X-ray emission flux in the five peaks (assuming they are at the same distance), but it may be the foreground star 2MASS J23045696+1220390 if we consider the error in pointing (the mean error of 8$''$) and resolution (spatial resolution 6$''$ FWHM) of XMM-Newton telescope. Target X2 and X3 are both located in the west strong spiral arm. X2 can be detected obviously from the high-resolution optical image from HST. In the position X3, there is an optical counterpart with very weak emission, which is only $\\sim$6$''$ from the compact source H4, also not far from two faint sources H3 and H5. H1 and H2 are located on the strong arm and near the stellar bar, especially H2 can not be resolved even in the images of HST. The potential target L lies on the north of the galactic nucleus in the bar, where there is little X-ray emission but an bright optical counterpart. The detailed information is listed in Table~\\ref{tabl4}, the coordinates of X-ray candidates are from XMM-Newton detection, and their luminosity (L$_X$) are including the emission from 0.2keV to 12keV.  For further investigations of properties of candidates, we made aperture photometry with radius of 6$''$ to get their SEDs (Figure~\\ref{fig10}). Due to the low spatial resolution, we used the sum flux of X3 and H4 to plot their SED. We also obtain the VSG-to-PAH ratio using F(24$\\mu$m)/F(8$\\mu$m) and surface SFR using Eq.~\\ref{eq_SFR2} with small photometric apertures (radii of $\\sim$2$''$) (see Table~\\ref{tabl4}). We found 3 sources have the positive ratios and 4 sources have surface SFR higher than 0.3 M$_{\\odot} $yr$^{-1}$ kpc$^{-2}$. By comparing these plots, we found that sources L and H3 have similar SEDs with the nucleus C, and they also have higher VSG-to-PAH ratios, which indicates they are more likely in the criteria region of AGNs \\citep{Wu07}. Thus, under the consideration of the two aspects, we suggest that the sources L and H3 are more likely remnants of the minor merger event."
    },
    "1107/1107.0464_arXiv.txt": {
        "abstract": "{  \\textit{Context}. OH maser emission is known to be associated with high mass star forming regions. Towards some of these regions, OH masers are associated with HII regions. Towards others, believed to be in an earlier evolutionary state, OH masers are offset from HII regions. Towards these later regions, it is believed that OH masers are associated with the circumstellar disk (e.g. Edris et al. \\citealp{edris05}; Gray et al. \\citealp{gray03}). These disks should be hosting dense dust grains. The presence of the hot dust could be traced via the millimeter continuum emission as well as IR emission.  \\textit{Aims}. studying the association between millimeter (mm) continuum, the OH masers emission, and IRAS sources.  \\textit{Methods}. A sample of 27 High Mass Star Forming Regions (HMSFRs) chosen from IRAS catalog and show OH maser emission (Edris et. al. \\citealp{edris07}) have been studied at 1.1 millimeter (mm) continuum emission of the Bolocam Galactic Plane Survey (BGPS).  \\textit{Results}. The 1.1-mm continuum emission have been found within $\\simeq$ 30$'$ towards 23 sources of the OH maser sample. These sources were divided into three groups depends on the offset of the closest mm peak from the OH maser position. The association between the OH, mm and IR emissions types have been confirmed for two sources. Generally the IRAS position is more consistent with the mm peaks than the OH maser emission and towards 10 sources the IRAS and OH masers are not consistent with the same mm peak.  \\textit{Conclusion}. The relatively large positional uncertainty do not allow to firm conclusions but it seems that the IR peak is closer to the mm emission than the OH maser. ",
        "introduction": "Compact HII regions, molecular outflows and circumstellar disks are signs of the existence of massive protostars (e.g. Garay \\& Lizano \\citealp{garay99}; Churchwell \\citealp{churchwell02}). Maser emission is also found to be associated with these objects (e.g. Garay \\& Lizano \\citealp{garay99}; Edris et al. \\citealp{edris05}) with OH one of most widespread types of maser associated with these regions. OH masers are associated with two different stages of star formation. The first one is associated with circumstellar disks and molecular outflows (Cohen, Rowland \\& Blair \\citealp{cohen84}; Brebner \\citealp{brebner88}; Cohen et al. \\citealp{cohen03}; Edris et al. \\citealp{edris07}). the other one is a relatively advanced star forming stage of the appearance of UCHII region (e.g. Garay \\& Lizano \\citealp{garay99}, and references therein). The mm emission could trace different phenomena of star formation process. It could trace slowly collapsing rotating molecular fragment onto itself, a circumstellar disk and molecular outflows (e.g. Zinnecker et al. \\citealp{zinnecker92} ; Adams et al. 1990 ; Beuther et al \\citealp{beuther02b}). therefore the mm emission should be associated with the first type of OH masers mentioned above. The association between the maser emission and mm and sub-mm have studied in case of methanol masers by Breen et al. (2010) and water masers by Jenness et al. (1995). The association of OH masers and the mm emission has not been studied. With this in mind a sample of high mass protostellar objects has been studied at the mm emission. ",
        "conclusions": ""
    },
    "1107/1107.3889_arXiv.txt": {
        "abstract": "The hydrodynamic evolution of the common envelope phase of a low mass binary composed of a $1.05 \\msun$ red giant  and a $0.6 \\msun$ companion has been followed for five orbits of the system using a high resolution method in three spatial dimensions.  During the rapid inspiral phase, the interaction of the companion with the red giant's extended atmosphere causes about 25\\% of the common envelope to be ejected from the system, with mass continuing to be lost at the end of the simulation at a rate $\\sim 2 \\msun {\\rm\\ yr}^{-1}$. In the process the resulting loss of angular momentum and energy reduces the orbital separation by a factor of seven. After this inspiral phase the eccentricity of the orbit rapidly decreases with time.  The gravitational drag dominates hydrodynamic drag at all times in the evolution, and the commonly-used Bondi-Hoyle-Lyttleton prescription for estimating the accretion rate onto the companion significantly overestimates the true rate. On scales comparable to the orbital separation, the gas flow in the orbital plane in the vicinity of the two cores is subsonic with the gas nearly corotating with the red giant core and circulating about the red giant companion. On larger scales, 90\\% of the outflow is contained within 30 degrees of the orbital plane, and the spiral shocks in this material leave an imprint on the density and velocity structure. Of the energy released by the inspiral of the cores, only about 25\\% goes toward ejection of the envelope. ",
        "introduction": "\\medskip  The common-envelope (CE) phase of binary star evolution, a brief phase during which two stars spiral together within a differentially rotating envelope leading to the loss of mass and orbital angular momentum from the system, was first proposed as an important stage in the evolution of close binaries during the 1970s \\citep{paczynski_common_1976}. Because close binary systems are central to the interpretation and understanding of many astronomical phenomena, including Type Ia supernovae, cataclysmic variables, X-ray binaries, and possibly short-duration gamma-ray bursts, developing a physical understanding of the CE phase has become important for predicting the formation rates and characteristics of these systems \\citep{2000ARA&A..38..113T}. Mapping the portions of binary parameter space that lead to different CE outcomes is a particularly important goal.  A definitive understanding of the CE phase requires computationally intensive hydrodynamical simulations of which a number of recent studies have been carried out \\citep[e.g.,][]{ricker_interaction_2008,2010NewAR..54...65T,passy_ce}. However, the final state of the post-CE systems has yet to be determined from such calculations and, hence, modeling of its outcomes in binary population synthesis studies has relied upon simple analytical prescriptions based on energetic considerations \\citep{1985ApJS...58..661I} where an efficiency for the conversion of orbital energy of the binary to the kinetic energy of the outflow is assumed.  Alternatively, a prescription based on angular momentum considerations has been proposed \\citep{2000A&A...360.1011N,2005MNRAS.356..753N}, where the specific angular momentum of the ejected matter is taken as a parameter.  Recent studies have focused on improvements to the former analytical prescription by developing a more accurate description of the conditions required for ejection: that is, the binding energy of the envelope and the response of the stellar interior to mass loss \\citep{2010ApJ...716..114X,2010Ap&SS.329..243G,2010ApJ...719L..28D,2010arXiv1009.5400L}. However, the efficiency of the ejection process remains to be determined.  Toward the goal of understanding the physical mechanisms that are important in governing the ejection of material during the CE phase and ultimately the efficiency of the process, we have carried out a high-resolution adaptive mesh refinement (AMR) simulation of the evolution of a binary initially consisting of a 1.05$\\msun$ red giant containing a 0.36$\\msun$ degenerate core with a 0.6$\\msun$ companion. The use of AMR in this study represents a major improvement over previous work based on stationary nested grids \\citep{1998ApJ...500..909S}.  This method ensures that the deep interior of the common envelope is always well resolved as matter moves about the center of mass of the system. Hence, the method permits one to study systems with binary components of nearly equal masses, which could not be adequately modelled with the stationary nested grid technique.  In these systems, the CE phase is initiated by Roche lobe overflow of the red giant star rather than by a tidal (Darwinian) instability of the orbit.  The results from the first 41 days of our simulation were described in Ricker and Taam (\\citeyear{ricker_interaction_2008}, hereafter Paper~I), wherein we showed that the assumption of Bondi-Hoyle-Lyttleton \\citep[BHL; ][]{hoyle_effect_1939,bondi_mechanism_1944, bondi_spherically_1952} accretion onto the companion dramatically overestimates the true accretion rate since the conditions for BHL accretion (uniform supersonic motion with the gravitational wake trailing the motion of the accretor) are not satisfied.  In this paper we examine the intermediate time evolution of the simulation whose early stages were described in Paper~I. Here, we focus specific attention on the determination of the drag forces and the mechanisms for angular momentum transport in the system. To gain a deeper understanding of the influence of the interaction on the core of the red giant and the companion, we also study the character of the mass flow in the vicinity of each component and the accretion onto each component. In addition, we examine the angular and radial distribution of the matter ejected from the CE, which may provide observational evidence that a system has passed through this stage, independent of the outcome.  In \\S 2 we briefly discuss the numerical methods and initial conditions used in our setup. The results are presented in \\S 3, and we conclude with a discussion of their implications in the last section. ",
        "conclusions": "We have used a large AMR simulation to study the CE evolution of a binary system initially consisting of a 1.05$\\msun$ red giant with a 0.36$\\msun$ degenerate core and a 0.6$\\msun$ companion. The use of AMR allows us to simulate stellar pairs with mass ratios close to unity at high resolution. We have followed the system through 57 days of evolution, during which time it has evolved through five orbital revolutions. The orbital separation decreases by a factor of 7 to $8.6 R_\\odot$, and the eccentricity steadily decreases with time to a value of 0.08. The CE interaction has led to the unbinding of 0.18$\\msun$, or 26\\%, of the red giant's initial gaseous envelope at this point in time.  We note that this fraction of unbound mass is larger than in the work by \\cite{2000ApJ...533..984S} (typically less than 15\\%), but is consistent with the trend in which this unbound mass fraction increases with a larger mass ratio (companion to the red giant) of the system.  Based on the release of orbital energy, the efficiency of the mass ejection process is $\\sim 25\\%$. Given that the mass loss rate from the system is $\\sim 2 \\msun$ yr$^{-1}$, the remaining mass is expected to be removed within 2 additional months (comparable to the time scale of the orbital decay) provided that the mass loss rate remains approximately constant. This is a reasonable assumption because the binding energy of the portion of the envelope that remains bound is increasing linearly with time at the end of the simulation. Additional orbital decay will result in the ejection of the remaining envelope.  It is possible that ejection of the envelope will be incomplete as some matter may fall back, and the effect of any material surrounding the system in the form of a circumbinary disk \\citep{2011arXiv1105.5698K} may lead to further orbital decay.  It is likely that this system will survive the CE phase with an orbital period of less than 3 days.  As in our previous study concerning the early stages of the evolution of this model system, the gravitational drag acting on the two stellar cores due to the nonsymmetrical gas distribution dominates over the hydrodynamical drag. The flow pattern in the vicinity of the cores in the orbital plane at 57~days is nearly uniform for the red giant core and circular for the companion, and in both cases it is subsonic. Such a description of the flow renders the assumptions underlying simple estimates of the drag based on a BHL prescription suspect and should be noted when considering evolutionary calculations of lower dimensionality.  The commonly used BHL prescription for estimating the mass accretion rate onto the companion overestimates the actual accretion rate by nearly two orders of magnitude. As discussed in Paper~I, this can have important implications for the likelihood of the possibility for accretion-induced collapse of an inspiralling neutron star in a stellar envelope to a black hole.  We note, however, that for the less evolved red giant star that we have modelled in the present investigation, the rate that we infer from our calculations can still exceed the rate for hypercritical accretion \\citep[$\\sim 10^{-3} \\msun$ yr$^{-1}$;][]{1993ApJ...411L..33C} suggesting that the outcome of the accretion evolution of an inspiralling neutron star is dependent on the evolutionary state of its giant companion.  An alternative prescription for the CE drag and resulting inspiral has been proposed by Meyer and Meyer-Hofmeister (1979; MM)\\nocite{meyer_formation_1979}, who consider a model for the CE evolution of a 5~$M_\\odot$ red giant and a 1~$M_\\odot$ main-sequence star. In this model, the system is treated as an inner rigidly rotating binary system frictionally interacting with an outer envelope undergoing differential rotation. The angular momentum transfer and energy dissipation in the envelope derives from turbulent convection. Gravitational drag is neglected in the envelope. Because gravitational drag appears to be the dominant dissipation mechanism in our calculation, we would argue that at least during the rapid inspiral phase the MM model greatly underestimates the drag and hence overestimates the inspiral timescale (see their Figure~5). Moreover, the radially averaged specific entropy profile of our system at $t = 56.7$~days is convectively stable, so if convective turbulence should have developed but did not because of our resolution, its integral scale and hence effective mixing length should be smaller than our minimum zone spacing ($2\\times10^{10}$~cm). The mixing length in MM is of order the pressure scale height, which in our calculation is $\\sim 7\\times10^{10}$~cm at its smallest. Hence, the turbulent viscosity is correspondingly overestimated in their model. In our Figure~\\ref{Fig:density} it is clear that the velocity field in the central regions is laminar on these scales. Instead, in our calculation angular momentum is removed from the envelope interior regions by spiral shocks.  It is conceivable that physics not included in our simulation could help to invert the entropy gradient and thus drive turbulence on small scales, changing the nature of the angular momentum transport. Our equation of state includes radiation pressure, but the gas is treated as being fully ionized, so energy derived from recombination is not included. If significant amounts of recombination were to occur within the inner regions, it could both raise the entropy there and contribute to the ejection of more envelope material. However, a rough estimate of the optical depth due to electron scattering (which dominates at the temperatures present in the inner regions) shows that gas and radiation remain tightly coupled at least up to a radius of $5\\times 10^{12}$~cm. Thus recombination should be important only in the very outer regions of the envelope and not in the inner region where the gravitational binding is greater. In more evolved red giants, where more of the mass is located at larger radii, recombination may play a greater role in unbinding mass, though in such cases one would need to consider the efficiency with which recombination photons couple to the surrounding gas.  The existence of distinct shells of outflowing material corresponding to the different orbits suggests a means for detecting systems that are undergoing or have recently undergone the common envelope phase. Once the outflow expands sufficiently to become optically thin, any spectral lines emitted by the gas will display Doppler-broadened profiles containing discontinuous jumps or features corresponding to the velocity jumps seen in Figure~\\ref{Fig:spiral-shocks}. The highest-velocity components result from the tangential motion at radii corresponding to several times the orbital separation.  As the outflow expands further, a more radial flow at larger distances is expected. Such components would be dependent on the orbital phase of the system and evolve secularly with time. These jumps should be detectable by spectroscopic observations with resolving powers ($\\lambda/\\Delta\\lambda$) of a few thousand or better, easily within the capabilities of current technology. The primary difficulty with such observations would be the short time available before the outflow density drops below the detection threshold. Discovery of post-common-envelope systems using this technique would therefore require a spectroscopic monitoring program following a large number of stars, or spectroscopic followup of optical transients detected in a large-area photometric survey with a high cadence (such as LSST)."
    },
    "1107/1107.4521_arXiv.txt": {
        "abstract": "Using Monte Carlo simulations, we model the luminosity distribution of recycled pulsars in globular clusters as the brighter, observable part of an intrinsic distribution and find that the observed luminosities can be reproduced using either log-normal or power-law distributions as the underlying luminosity function. For both distributions, a wide range of model parameters provide an acceptable match to the observed sample, with the log-normal function providing statistically better agreement in general than the power-law models. Moreover, the power-law models predict a parent population size that is a factor of between two and ten times higher than for the log-normal models. We note that the log-normal luminosity distribution found for the normal pulsar population by Faucher-Gigu\\`ere and Kaspi is consistent with the observed luminosities of globular cluster pulsars. For Terzan~5, our simulations show that the sample of detectable radio pulsars, and the diffuse radio flux measurement, can be explained using the log-normal luminosity law with a parent population of $\\sim 150$ pulsars. Measurements of diffuse gamma-ray fluxes for several clusters can be explained by both power-law and log-normal models, with the log-normal distributions again providing a better match in general. In contrast to previous studies, we do not find any strong evidence for a correlation between the number of pulsars inferred in globular clusters and globular cluster parameters including metallicity and stellar encounter rate. ",
        "introduction": "\\label{lb:introductions}  The first millisecond radio pulsar in a globular cluster (GC) was discovered by \\citet{lbmkb87}, shortly after earlier predictions that the putative progenitors of millisecond pulsars (MSP) are low-mass X-ray binaries \\citep{acrs82} which are known to be present in GCs \\citep{katz75}. Inspired by this discovery, a large number of sensitive pulsar searches have been performed over the years resulting in the currently observed population of 143 radio pulsars in 27 GCs\\footnote{For a complete list, see http://www.naic.edu/$\\sim$pfreire/GCpsr.html.}. GCs have some physical properties which are different than those of the galactic disk. Examples are extremely high stellar density and high abundance of metal poor population II stars indicating that they were born in the early phase of the Galaxy's formation. These facts lead naturally to the question whether the population of radio pulsars in GCs is different to their counterparts in the Galactic disk. A number of differences are already well-known. In particular, the high abundance of MSPs both in eccentric binary systems and as isolated objects. These phenomena can be explained as the results of two-body or three-body stellar interactions in the dense stellar environment of GCs \\citep{rkb87, vhpr87, ihrbf08, br09}. On the other hand, the luminosity of a pulsar is a more fundamental property as it can in principle be linked to the pulsar emission mechanism. It is therefore important to establish whether there is any evidence for a different luminosity function of MSPs in the disk compared to those in GCs.  There are two main ways to determine the pulsar luminosity function numerically: (i) a full dynamical approach; (ii) a snapshot approach. In the dynamical approach, a simulation is performed in which a model galaxy of pulsars is seeded according to various prescriptions of birth locations and initial rotational parameters. Each of these synthetic pulsars is then ``evolved'' both kinematically in a model for the Galactic gravitational potential and rotationally using a model for neutron star spin-down. The resulting population is then passed through the various detection criteria \\citep[see, e.g.,][]{fk06}. The resulting set of ``detectable'' pulsars is then compared to the observed sample. The snapshot approach differs from the dynamical approach in that the pulsars are seeded at their final positions in the galaxy without assuming anything about their spin-down or kinematic evolution and thus form a picture of the present day population.  The dynamical approach has been used extensively to study normal pulsars in the galactic disk \\citep{bwhv92, goboh02, fk06, rl10}. These studies considered the luminosity of pulsars to be described by power-law functions involving $P$ and $\\dot{P}$ with a substantial dispersion to account for distance uncertainties and beam geometry. One of the conclusions by \\citet[][hereafter FK06]{fk06}, also verified by \\citet{rl10} is that the resulting parent population of luminosities appears to be well described by a log-normal function. As discussed further in Section 3.1, the log-normal parameters favored by FK06 (for the base-10 logarithm of the 1400-MHz luminosities) are a mean of --1.1 and a standard deviation of 0.9. One of the goals of the current study is to examine whether these log-normal parameters are consistent with the observed populations of GC pulsars.  For pulsars in GCs, where it is difficult to model the effects of stellar encounters and the cluster potential, the dynamical approach has so far not been used for radio pulsar population syntheses. Although such an approach may be tractable in future, we will adopt in this work a version of the snapshot approach. We will carry out Monte Carlo simulations that assume all GCs have the same intrinsic pulsar luminosity function, but a different population size and use this approach to model the observed sample of pulsars given the various ranges of luminosities. As we shall see, this approach provides a remarkably good agreement between the model and observed luminosity distributions.  Within the snapshot framework, one usually fits the complementary cumulative distribution function (CCDF) of pulsar luminosity as a power law as $N ( > L_{1400}) = N_0 L_{1400}^q$, where $L_{1400}$ is the luminosity of the pulsar at 1400 MHz, $N$ is the number of pulsars having luminosity value greater than $L_{1400}$, $N_0$ and $q$ are constants. We note here that sometimes in literature \\citep{hrskf07, hct10}, the CCDF has been mentioned as the cumulative distribution function; but actually the cumulative distribution function (CDF, $N ( \\leq L_{1400})$ ) is related to the CCDF as CDF = $1-$~CCDF. Using this snapshot technique, \\citet{hct10} concluded that the luminosities of MSPs in GCs are different from those in the galactic disk as the CCDF for GC MSPs is much steeper than that of disk pulsars. This is a very important conclusion. If correct, it would imply that the radio luminosity is related to differences in formation processes between the disk and GC pulsars. The same analysis was re-performed with more recent distance estimates of GCs and the resultant CCDF was even steeper \\citep{bl10}.  Since GCs are generally at large distances, the luminosity function of observed pulsars is not as well sampled as in the Galactic disk.  In the present work we try to account for this incompleteness by considering GC MSPs as the brighter tail of some intrinsic parent population. The goal of the current work is to explore the range of possible distributions that are consistent with the current sample of GC pulsars. The plan for the rest of this paper is as follows. In Section \\ref{sec:data}, we describe the pulsar sample we use. In Section \\ref{sec:anal}, we present our analysis procedure and its main results. In Section \\ref{sec:constraints}, we investigate additional constraints on allowed model parameters from observations of diffuse radio and gamma-ray flux. In Section \\ref{sec:comp}, we compare our results with earlier work. We draw our main conclusions in Section \\ref{sec:conclu}. ",
        "conclusions": ""
    },
    "1107/1107.3137_arXiv.txt": {
        "abstract": "We present the results of a 42-orbit {\\it Hubble Space Telescope} Wide-Field Camera 3 ({\\it HST}/WFC3) survey of the rest-frame optical morphologies of  star forming galaxies with spectroscopic redshifts in the range $z=1.5-3.6$.  The survey consists of 42 orbits of F160W imaging covering $\\sim 65$ arcmin$^2$ distributed widely across the sky and reaching a depth of  27.9 AB for a $5\\sigma$ detection within a 0.2 arcsec radius aperture. Focusing on an optically selected sample of 306 star forming galaxies with stellar masses in the range $M_{\\ast} = 10^9 - 10^{11} M_{\\odot}$, we find that typical circularized effective half-light radii range from $\\sim 0.7 - 3.0$ kpc and describe a stellar mass - radius relation as early as $z\\sim3$. While these galaxies are best described by an exponential surface brightness profile (Sersic index $n \\sim 1$), their distribution of axis ratios  is strongly inconsistent with a population of inclined exponential disks and is better reproduced by triaxial stellar systems with minor/major and intermediate/major axis ratios $\\sim 0.3$ and 0.7 respectively. While rest-UV and rest-optical morphologies are generally similar for a subset of galaxies with {\\it HST}/ACS imaging data, differences are more pronounced at higher masses $M_{\\ast} > 3 \\times 10^{10} M_{\\odot}$. Finally, we discuss galaxy morphology in the context of efforts to constrain the merger fraction, finding that morphologically-identified mergers/non-mergers generally have insignificant differences in terms of physical observables such as stellar mass and star formation rate, although merger-like galaxies selected according to some criteria have statistically smaller effective radii and correspondingly larger $\\Sigma_{\\rm SFR}$. ",
        "introduction": "In recent years our understanding of the broad global characteristics of galaxies in the young universe has grown considerably.  Using rest-frame UV and optical spectroscopy and multi-wavelength broadband photometry it has been possible to estimate their stellar and dynamical masses, average metallicities, ages, and star formation rates across cosmic time from $z>6$ to the present day (e.g., Shapley et al. 2005; Cowie \\& Barger 2008; Maiolino et al. 2008; Stark et al. 2009).   Such studies indicate that the majority of structures observed in the local universe were already in place at $z\\sim1$ (Papovich et al. 2005) and point to the era spanned by the redshift range $1.5<z<3$ as the peak epoch of both the cosmic star formation rate density (Dickinson et al. 2003; Reddy et al. 2008) and AGN activity in the universe (e.g., Miyaji et al. 2000).  In contrast to our knowledge of the global characteristics of such galaxies from ever-expanding samples however, our understanding of their internal structure and evolution has been limited by their small angular size. It has therefore been challenging to constrain the major mode of mass assembly in these galaxies (i.e., from major/minor mergers, hot mode or cold filamentary gas accretion, etc.). With typical half-light radii $\\sim 0.2-0.3$ arcsec at $z\\sim2$ (Bouwens et al. 2004; Nagy et al. 2011), such galaxies are barely resolved in the $\\sim1$ arcsec FWHM ground-based imaging and spectroscopy that form the backbone of the observational data.   Significant efforts have therefore been invested in imaging studies capitalizing on the high angular resolution afforded by the {\\it Hubble Space Telescope} ({\\it HST}). Early efforts to characterize the morphologies of galaxies at $z\\sim1.5-3$ (e.g., Abraham et al. 1996; Giavalisco et al. 1996; Lowenthal et al. 1997; Bouwens et al. 2004; Conselice et al. 2004; Lotz et al. 2006; Papovich et al. 2005; Law et al. 2007b; and references therein) used the visible-wavelength surveying efficiency of the ACS (Advanced Camera for Surveys) to demonstrate that star forming galaxies typically have irregular, clumpy morphologies unlike the well-known Hubble sequence that has been established since $z\\sim1$ (e.g., Conselice et al. 2005, Oesch et al. 2010). Indeed, % rest-UV luminosity and morphology for such galaxies appears to be only poorly correlated with other physical observables such as stellar mass, outflow characteristics, and characteristic rotation velocity (e.g., Shapley et al. 2005; Law et al. 2007b).  Recent technological developments have permitted additional insights to be gleaned from ground-based observations using adaptive-optics (AO) fed imagers or integral field unit (IFU) spectrographs on 10m-class telescopes (e.g., Law et al. 2007a, 2009; Melbourne et al., 2008ab, 2011; Stark et al., 2008; F{\\\"o}rster-Schreiber et al. 2009; Wright et al. 2009; Jones et al., 2010). Such IFU spectroscopy mapping rest-frame optical nebular line emission  (redshifted into the near-IR at $z>1$) from star forming galaxies has suggested that high redshift star forming galaxies often have dispersion-dominated kinematics at odds with the classical picture of galaxy formation via rotationally supported thin gas disks (e.g., F{\\\"o}rster-Schreiber et al. 2009; Law et al. 2009). Instead, the dynamical evolution of these systems may be driven by gravitational instabilities within massive gas-rich clumps or low angular-momentum cosmological gas flows (e.g., Kere{\\v s} et al. 2005; Bournaud et al. 2007; Genzel et al. 2008 ). What is immediately clear is that we do not yet understand the dynamical state of galaxies during the period when they are forming the majority of their stars.   Both early rest-UV imaging and AO IFU observations of high-redshift galaxies tend to trace regions of active star formation however, and in order to understand these galaxies we also wish to map the regions in which the bulk of the underlying stellar population live. While young and old populations may have a generally similar distribution for lower-mass galaxies (e.g., Conselice et al. 2011), more significant differences exist for galaxies with larger stellar mass (e.g., Dickinson 2000; Papovich et al. 2005; and \\S \\ref{uvopt.sec}). Efforts to characterize rest-optical galaxy morphologies using ground-based instruments and/or the {\\it HST}/NICMOS camera have been made by (e.g.) Papovich et al. (2005), Franx et al. (2008), Toft et al. (2009), van Dokkum et al. (2010), and Mosleh et al. (2011), generally finding that galaxies at $z\\sim2$ were significantly smaller at fixed stellar mass than in the local universe. Additionally,  {\\it HST}/NICMOS work by Kriek et al. (2009) has demonstrated that star-forming and quiescent galaxies differ substantially from each other in relative compactness of their rest-optical morphologies, and both differ from their kin in the local universe.  Given the narrow field of view of both ground-based AO-fed imagers (e.g., Carrasco et al. 2010) and the {\\it HST}/NICMOS camera (e.g., Conselice et al. 2011a) however, it is only recently with the advent of the new WFC3 camera onboard {\\it HST} that it has become practical to perform wide-field morphological surveys in the near-IR that trace rest-frame optical emission from galaxies at $z\\gtrsim1$. The results of the first such studies in the UDF have been reported recently in the literature (e.g., Cameron et al. 2010; Cassata et al. 2010; Conselice et al. 2011b). Our recent survey has greatly extended these early results by obtaining {\\it HST}/WFC3-IR morphological data for 306 $z = 1.5-3.6$ galaxies in 10 fields widely distributed across the sky for which we have obtained dense spectroscopic sampling.  Preliminary results for the evolution of the stellar mass - radius  relation were presented in Nagy et al. (2011). In this first contribution of a series of papers using the full sample, we introduce our survey and describe a selection of results concerning  evolution of the characteristic size, shape, and major merger fraction for actively star forming galaxies. Future contributions (Law et al. 2012, in preparation) will discuss the relation between morphology and low-ionization gas-phase kinematics, treat quiescent galaxies and AGN, and discuss the morphology of uniquely interesting galaxies (e.g., Q2343-BX442; Law \\& Shapley 2012, in preparation) in greater detail. This paper is organized as follows:  In \\S \\ref{obs.sec} we describe the {\\it HST}/WFC3 observing program and review the properties of the star forming galaxy sample. In \\S \\ref{stats.sec} we present postage-stamp morphologies of the galaxy sample and discuss our morphological analysis techniques. An extended discussion of the robustness of the morphological statistics and the systematic variations between measurement systems commonly adopted in the literature is presented in the Appendix. \\S \\ref{basiccharac.sec} summarizes the basic morphological characteristics (luminosity profile, relation to rest-UV imaging, and intrinsic 3D shape) of the galaxy sample, and the implications of our data for the evolution of the stellar mass - effective radius relation are discussed in \\S \\ref{massrad.sec}. Finally, we use a variety of morphological statistics to constrain the major merger fraction and its evolution with redshift in \\S \\ref{mergers.sec}. We summarize our results in \\S \\ref{discussion.sec}.  Throughout our analysis, we adopt a standard $\\Lambda$CDM cosmology based on the seven-year WMAP results (Komatsu et al. 2011) in which $H_0=70.4$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.272$, and $\\Omega_{\\Lambda}=0.728$. ",
        "conclusions": "\\label{discussion.sec}  We have presented rest-optical morphologies for a sample of 306 spectroscopically confirmed  $z = 1.5-3.6$ star forming galaxies with stellar masses in the range $M_{\\ast} = 10^9 - 10^{11} M_{\\odot}$. Since these galaxies were distributed among 10 different fields widely separated on the sky the effects of sample variance are expected to be greatly reduced compared to surveys over contiguous regions of similar total area. We summarize our principle scientific conclusions as follows:  \\begin{enumerate}  \\item Typical $z \\sim 1.5-3.6$ star forming galaxies have circularized effective radii $r_e \\approx 0.7-3$ kpc and a projected $n\\sim1$ exponential surface brightness profile that extends out to $>6 r_e$ in stacked galaxy images. The observed sizes are consistent with previous observational estimates (e.g., Buitrago et al. 2008; Kriek et al. 2009) for high-mass galaxy populations and with numerical simulations (e.g., Sales et al. 2010) that assume strong stellar feedback.  \\item A stellar mass - radius relation for star forming galaxies is observed to exist as early as $z\\sim3$; at fixed mass typical sizes evolve with redshift as $\\sim (1+z)^{-1.07 \\pm0.28}$ in the interval $z \\sim 3$ to $z\\sim 1.5$. These galaxies must grow at least as fast as $r \\sim M_{\\ast}$ in order  to evolve onto the local late-type galaxy relation by the present day.  \\item The distribution of axis ratios $b/a$ is strongly inconsistent with a population of axisymmetric thick exponential disks and more consistent with a population of triaxial ellipsoids with intrinsic minor/major and intermediate/major axis ratios $0.3\\pm0.2$ and $0.7\\pm0.1$ respectively. The typical ellipticity is qualitatively similar to that previously found by Ravindranath et al. (2006), but there may be mild evidence for evolution with wavelength. The ellipsoidal nature of these galaxies indicates at minimum that the distribution of stellar mass within them is markedly asymmetric, and (in combination with their high gas fractions and velocity dispersions) may further suggest that they are not in stable dynamical equilibrium with short-lived gas disks (e.g., Ceverino et al. 2010) continually forming and re-forming from recently accreted gas until stabilized (e.g., Martig et al. 2010) by a sufficiently massive triaxial stellar component.   \\item Consistent with previous studies (e.g., Dickinson 2000; Papovich et al. 2005), rest-optical ($\\lambda \\sim 4000-5000$\\AA) and rest-UV ($\\lambda \\sim 2000-3000$ \\AA) morphology for $z\\sim2$ star forming galaxies is generally similar with typical color dispersion $\\xi \\sim 0.02$ (although rest-UV radii are larger by $21\\pm2$\\% on average), while high mass ($M_{\\ast} > 3 \\times 10^{10} M_{\\odot}$) galaxies tend to exhibit greater morphological differences with $\\xi$ as large as $0.28$ (although c.f. Bond et al. 2011). The most massive galaxies in our sample are typically bright and well nucleated at rest-optical wavelengths but faint and diffuse in the rest-UV.    \\item Finally, we demonstrate that while the nonparametric morphological statistics $G$, $M_{20}$, $C$, $A$, and $\\Psi$ calculated using different segmentation maps commonly adopted in the literature are strongly correlated with each other, there can be systematic offsets that are important to account for when comparing values between samples or estimating merger fractions.  Merger fractions estimated according to the $G-M_{20}$, $A$, or pair criteria are consistent with recent determinations in the literature (e.g., Conselice et al. 2011a), with evidence for at most mild evolution with redshift.  There is moderate overlap between galaxies selected as mergers with these three criteria, but in general mergers and non-mergers have statistically indistinguishable distributions of measured and inferred properties (SFR, stellar mass, etc.), with the exception that mergers selected by the pair and $G-M_{20}$ statistics have smaller effective radii and correspondingly larger $\\Sigma_{\\rm SFR}$. We suggest that  most  $z\\sim 2-3$ star forming galaxies may be dynamically unstable systems driven by the accretion of large quantities of gas, whether this gas is acquired through mergers, cold-mode, or hot-mode accretion processes.  \\end{enumerate}   In general, our observations are consistent with inside-out growth of star forming galaxies in the young universe.  We suggest that mass growth proceeds according to the following qualitative picture: `Typical' $z\\sim2$ star forming galaxies appear to be gas-rich, compact, triaxial systems systems that are dominated by velocity dispersion between individual star forming regions rather than systemic rotation, and whose high $\\Sigma_{\\rm SFR}$ drives strong outflows into the surrounding IGM.  As these galaxies mature they gain stellar mass, stabilizing the formation of extended (albeit still thick) gaseous disks in which rotational support plays an increasing role.  As the star formation migrates from central regions into these extended disks the $\\Sigma_{\\rm SFR}$ drops, and the disk component superimposes a zero-velocity component atop the outflowing absorption line gas (Law et al. 2011, in preparation)."
    },
    "1107/1107.3140.txt": {
        "abstract": "As part of an extensive study of the physical properties of active galactic nuclei (AGN) we report high spatial resolution near-IR integral-field spectroscopy of the narrow-line region (NLR) and coronal-line region (CLR) of seven Seyfert galaxies. These measurements elucidate for the first time the two-dimensional spatial distribution and kinematics of the recombination line Br$\\gamma$ and high-ionization lines [Si~{\\sc vi}], [Al~{\\sc ix}] and [Ca~{\\sc viii}] on scales $<300$ pc from the AGN. The observations reveal kinematic signatures of rotation and outflow in the NLR and CLR. The spatially resolved kinematics can be modeled as a combination of an outflow bicone and a rotating disk coincident with the molecular gas. High-excitation emission is seen in both components, suggesting it is leaking out of a clumpy torus. While NGC 1068 (Seyfert 2) is viewed nearly edge-on, intermediate-type Seyferts are viewed at intermediate angles, consistent with unified schemes. %While the outflow axis in the Seyfert 2 galaxy NGC 1068 is almost perpendicular to %the line of sight, intermediate-type Seyferts are viewed at intermediate angles, %consistent with unified schemes. A correlation between the outflow velocity and the molecular gas mass in $r<30$ pc indicates that the accumulation of gas around the AGN increases the collimation and velocity of the outflow. %a massive torus produces fast and collimated outflows. The outflow rate is 2--3 orders of magnitude greater than the accretion rate, implying that the outflow is mass loaded by the surrounding interstellar medium (ISM). In half of the observed AGN the kinetic power of the outflow is of the order of the power required by two-stage feedback models to be thermally coupled to the ISM and match the $M_{BH}-\\sigma^*$ relation. In these objects the radio jet is clearly interacting with the ISM, indicative of a link between jet power and outflow power.  %In these objects a clear interaction of the ISM with the radio jet is observed, %indicative of a link between jet power and outflow power.  %AGN with greater amount of gas in their centers have collimated and fast %outflows. %This is interpreted as evidence for a collimating torus, which influences %considerably the properties of the outflows. %The accumulation of gas clouds towards the nucleus (higher gas mass) increments %the collimation power of the torus and produces outflows with higher velocities %(stops the outflow expanding sideways). %The outflow rate is 2--3 orders of magnitude greater than the accretion rate, %implying that the outflow is mass loaded by the surrounding interstellar medium %(ISM) in a two-stage process. %which indicates that the outflowing gas is not likely a single wind from the AGN, %but the result of a two-stage process in which the circumnuclear gas has been %pushed away by an already accelerated AGN wind and the radio jet.  %As part of an extensive study of the physical properties of active galactic nuclei %(AGN) we report here high spatial resolution near-IR integral field spectroscopy %of the narrow line region (NLR) and the coronal line region (CLR) of seven Seyfert %galaxies. These measurements elucidate for the first time the two-dimensional %spatial distribution and kinematics of the near-IR recombination line Br$\\gamma$ %and high ionization lines [Si~{\\sc vi}], [Al~{\\sc ix}] and [Ca~{\\sc viii}] on a %scale of less than a few parsecs from the AGN. We find that the CLRs are in %general more compact than the NLRs with total sizes varying from 160 to 300 pc, %and that the morphologies of the two regions at these scales are similar, %indicating that they are being produced by the same ionizing source. The %integral-field observations reveal kinematic signatures of rotation and outflow in %the NLR and CLR of Seyfert galaxies. We confirm previous claims of acceleration %along the NLR of NGC~1068 and NGC 4151, and discover that this is a common feature %in the kinematics of the CLR. Kinematic models incorporating these two types of %motion provide a good fit to the overall velocity fields. %Regarding the rotational component, the data clearly shows that the disks ionized %gas are spatially coincident with the galactic disks of molecular gas, suggesting %that the two types of gas are spatially mixed. Furthermore, the ionized gas %emission is observed also along its equatorial plane, indicating that the AGN %ionizes gas beyond the borders of the bicone.  %Regarding the outflowing component, a correlation between the outflow velocity %($v_{\\mathrm{max}}$) and the molecular gas mass ($M_{\\mathrm{gas}}$) in $r<30$ pc %has been found, in the sense that galaxies with greater amount of gas in their %centers have higher degree of collimation and higher outflow velocities. This is %interpreted as evidence for a geometrically thick collimating structure of %molecular gas (the torus invoked by unified models), which influences considerably %the properties of the outflows. The accumulation of gas clouds towards the nucleus %(higher $M_{\\mathrm{gas}}$) increments the collimation power of the torus and %produces outflows with higher velocities (stops the outflow expanding sideways). %Furthermore, the observed bipolar outflows are consistent with the torus model of %AGN if a bicone with internal and external faces is considered, i.e. it has a %thickness. %We have estimated mass outflow rates in the range 1 to 10 M$_\\sun$ yr$^{-1}$, %which are $\\sim10^2-10^3$ times the accretion rate necessary to feed the AGN. %The kinetic powers of the outflows ($\\dot{E}_{\\mathrm{out}}$) are $10^2-10^4$ %times smaller than the AGN bolometric luminosities. %We note that $\\dot{E}_{\\mathrm{out}}$ in half of the galaxies presenting outflows %is of the order of the power required by two-stage feedback models to be thermally %coupled to the circumnuclear gas ($\\dot{E}_{\\mathrm{out}}\\approx %0.005L_{\\mathrm{bol}}$), and match the $M_{BH}-\\sigma^*$ relation. These are also %the objects in which a clear interaction of the ISM with the radio jet is %observed, indicative of a link between jet power and $\\dot{E}_{\\mathrm{out}}$. The %high velocity components detected at the nucleus of some galaxies and the high %$\\dot{M}_{\\mathrm{out}}/\\dot{M}_{\\mathrm{acc}}$ ratios indicate that the %outflowing gas is not likely a single wind from the AGN, but the result of a %two-stage process in which the circumnuclear gas has been pushed away by an %already accelerated AGN wind and the radio jet. ",
        "introduction": "\\label{int}  The narrow-line region (NLR) in active galactic nuclei (AGN) plays a crucial role in attempts to understand the AGN structure and evolution, as well as the interaction between the nuclear engine and the circumnuclear interstellar medium (ISM) of the host galaxies. Among the fundamental AGN components, the NLR is the largest observable structure in the UV, optical and near-IR that is directly affected by the ionizing radiation and dynamical forces from the active nucleus. %being directly accessible via imaging and spatially resolved spectroscopy, at %least for nearby AGNs. The size and properties of the NLR has been scrutinized intensely over the past few decades. Ground-based and $Hubble$ $Space$ $Telescope$ ($HST$) [O~{\\sc iii}] $\\lambda5007$\\r{A} (hereafter [O~{\\sc iii}]) and H$\\alpha$ $+$ [N~{\\sc ii}] narrow-band imaging of Seyfert galaxies (Pogge 1988a, 1988b; Haniff et al. 1988; Mulchaey et al. 1996; Evans et al. 1991, 1993; Schmitt et al. 1996, 2003) revealed extended emission in several of these objects. The overall morphology of the NLR is varied, but there is a trend in the sense that the NLRs in Seyfert 2 galaxies are more elongated than those in Seyfert 1 galaxies. The fact that several Seyfert 2s presented conically shaped NLRs (e.g., NGC 1068 and NGC 4388) indicates that the source of radiation ionizing the gas is collimated, a result that gave strong support to the torus version of the Unified Model of AGN (Antonucci 1993; Urry \\& Padovani 1995). However, when comparing ground-based  [O~{\\sc iii}] images of Seyfert galaxies from Mulchaey et al. (1996) with the $HST$ snapshot survey of Schmitt et al. (2003), the latter reveal, on average, four times smaller NLR sizes, probably due to low sensitivity (quick exposures). Furthermore, one needs to be cautious when talking about ``NLR size'' as  [O~{\\sc iii}], or other low-ionization lines, can be produced by other physical processes not related to the AGN such as star formation or shock-ionized gas. %which postulates that a geometrically and optically thick torus outside of the BLR %produces a bicone of radiation that ionizes the NLR and that Seyfert type depends %on viewing angle, with type 2 objects having a bicone axis perpendicular to the %line-of-sight (LOS) and therefore conical (or biconical)-shaped NLRs, and type 1 %objects having a bicone axis nearly paralell to the LOS resulting in more compact %and circular NLRs.  %It is now widely accepted that the high-energy output of an active galactic %nucleus (AGN) is due to the accretion of material onto a supermassive black hole %(SMBH). The accretion process releases energy in the form of high-energy photons %which ionize %The central engine ionizes %the surrounding interstellar medium (ISM) of the host galaxy leading to the %broad-line region (BLR) in the inner part (tens of light-days, Peterson et al. %2004) and the narrow-line region (NLR) further out (several hundred parsecs, %Schmitt et al. 2003a, 2003b). %By definition, the NLR is responsible for emission lines with %widths $<400$ km s$^{-1}$ full width at half-maximum (FWHM), whereas emission %lines from the BLR have widths in the range $1000-5000$ km s$^{-1}$ FWHM. Also by %definition, type 2 AGNs show only narrow emission lines in their UV and optical %spectra, whereas type 1 objects show both broad and narrow emission lines. %According to the Unified Model of AGN, orientation effects play a major role in %explaining the observational appearance of AGNs. In the torus version of this %model a geometrically and optically thick torus of molecular gas and dust would %block our view of the BLR. In the case of type 2 AGNs the line-of-sight (LOS) %intercepts the torus (edge-on) and obscures the broad emission lines. In type 1 %the torus is orientated such that the LOS is perpendicular to the plane of the %torus (pole-on) allowing observations of the BLR. %Some of the most important evidence for this model is the %observation of a deficit of ionizing photons in Seyfert 2 galaxies %(Neugebauer et al. 1980; Wilson, Ward, \\& Haniff %1988; Storchi-Bergmann, Wilson, & Baldwin 1992) and the %collimated escape of radiation from the nucleus. %The NLR is the largest observable structure in the UV and optical that is directly %affected by the ionizing radiation and dynamical forces from the active nucleus. %Hence, studies of its physical conditions and kinematics should yield important %insights into the machinery of AGNs as %material is processed into and out of the environment surrounding %the central SMBH. %Therefore, the NLR in nearby AGN is an important laboratory to study in detail the %co-evolution of SMBHs and their host galaxies.  %The accretion process releases energy in the form of high-energy photons which %ionize material orbiting close to the SMBH, producing the Doppler-broadened %emission lines characteristic of the broad line region (BLR). %Moving outward further, there exists lower density gas which is known as the %narrow line region (NLR). In a virial sense, the greater distance from the SMBH %results in a lower velocity, leading to the narrow widths of the emission lines in %this region.  %Seyfert galaxies are important laboratories to study in detail the co-evolution of %supermassive black holes (SMBHs) and their host galaxies. The reason is simply %that these are the nearest objects in which the NLRs are directly accessible via %imaging and spatially resolved spectroscopy. %AGNs display many energetic phenomena (broad emission lines, narrow permitted and %forbidden optical emission lines, X-rays, relativistic jets, radio lobes) %originating from matter accreting onto a SMBH.  %Hence, studies of its physical conditions and kinematics should yield important %insights into the machinery of AGNs as material is processed into and out of the %environment surrounding the SMBH.  %The $Hubble$ $Space$ $Telescope$ ($HST$) has provided emission-line images of the %NLR in many Seyfert galaxies at a spatial resolution of $0.1\\arcsec$. The images %show that NLRs tend to be clumpy, with many resolved ``knots'' of emission that %are typically a few tenths of an arcsecond in size (Evans et al. 1991, 1993; %Schmitt et al. 2003b). The overall morphology of the NLR is varied; triangular, %linear, spiral (or S-shaped), elliptical, circular, arclike, and amorphous shapes %are all common. However, there is a trend in the sense that the NLRs in Seyfert 2 %galaxies are more elongated than those in Seyfert 1 galaxies, confirming previous %ground-based results.  %to the fact that in most of these galaxies the emission is concentrated %in the inner $2\\arcsec$ region and $HST$ is not sensitive enough to detect %halo-like [O iii] emission. These considerations question the definition of the %``NLR size'' from [Oiii] imaging alone.   %Ground-based long-slit spectroscopic studies have been used to provide the missing %kinematics and spectral information (e.g. Storchi-Bergmann, Wiklson & Baldwin %1992; Storchi-Bergmann et al. 1999; Cooke et al. 2000; Fraquelli, %Storchi-Bergmann & Binette 2000; Ruiz et al. 2001; Gonz\u00b4alez %Delgado et al. 2002; Fraquelli, Storchi-Bergmann & Levenson %2003). These studies  The physical conditions in the NLR have also been studied extensively, and it is now widely accepted that photoionization by the central continuum source plays a dominant role in exciting the gas (e.g. Kraemer \\& Crenshaw 2000a; Kraemer et al. 2000; Bennert et al. 2006) although shocks may play an important role in localized regions (Kraemer \\& Crenshaw 2000b; Dopita et al. 2002; Contini et al. 2009).  Ground-based long-slit spectroscopic studies attempted to determine the kinematics of the NLR, but a general consensus on the velocity flow pattern was not reached; cases were made for infall, rotation, outflow, etc. (e.g. Veilleux 1991; Moore \\& Cohen 1996; Cooke et al. 2000; Fraquelli et al. 2000; Ruiz et al. 2001; Fraquelli et al. 2003). %Since the NLR extends out to only a few arcsecs in most Seyfert galaxies (e.g. at %z=0.01 and H$_0=75$ km s$^{-1}$ Mpc$^{-1}$, 100 pc subtends only %$\\sim0.5\\arcsec$), these studies had to rely on spatially integrated line %profiles. %Unfortunately, similar profile shapes can be generated from a wide variety of %kinematic models. %Thus, spatially resolved spectra across %the NLR are essential for further kinematic constraints. %Early attempts with $HST$ made use of the long-slit spectroscopy mode %of the Faint Object Camera (FOC), resulting in kinematic information %on the NLRs in NGC~4151 (Winge et al. 1997, 1999), %NGC~1068 (Axon et al. 1998), and Mrk~3 (Capetti et al. 1999). Spectroscopic studies with the Faint Object Camera (FOC) on board of $HST$ also fail in this regard. Several models were advocated on the basis of these data, including rotation, radial acceleration by radio jets, and lateral expansion of gas around the jets (Winge et al. 1997, 1999; Axon et al. 1998; Capetti et al. 1999). %Despite being a well-studied galaxy, the %dynamics and excitation of the NLR, as well as the role of the %radio jet, are not yet fully understood. %We have found that that the kinematics of the NLR are dominated by radial outflow %in the four nearby Seyferts that we have studied to date: Observations with the Space Telescope Imaging Spectrograph (STIS) on $HST$ of a number of key objects have suggested that radial outflow is the dominant component of motion in the NLR. The first reliable evidence for outflow in the NLR came from $HST$/STIS long slit observations of NGC 1068 (Crenshaw \\& Kraemer 2000) and  NGC 4151 (Hutchings et al. 1998; Crenshaw et al. 2000). Since then, detailed constraints on kinematic models of the NLR in Seyfert galaxies became possible (see, e.g., Crenshaw et al. 2010 and references therein). %The objects studied include: NGC 4151 (Hutchings et al. 1998; Crenshaw et al. %2000; Das et al. 2005), NGC 1068 (Crenshaw \\& Kraemer 2000, Das et al. 2006, %2007), Mrk 3 (Ruiz et al. 2001), Mrk 573 (T. Fischer et al. 2010) and Mrk 78 %(Fischer et al. 2011). Although the comparison of the radial velocity curves with biconical models provide a good fit to the overall flow pattern, there exist significant local variations that remain unexplained. Furthermore, the dynamics, origin and energetics of the outflows, as well as the role of the radio jet and the presence of other types of motion in the NLR kinematics, are not fully understood. %possible kinematic components in the NLR, %Furthermore, the velocity curves in some objects, e.g. NGC~4151, can be well %matched with other kinematic models, as demonstrated by \\citet{winge99} who found %that, despite slight perturbations introduced probably by the radio jet, the %general [O~{\\sc iii}] velocity curve in NGC~4151 could still be well represented %by planar rotation. Two-dimensional velocity fields with high spatial resolution are needed to get a more complete picture of the kinematics in the NLR. Recent studies based on integral-field spectroscopy (a few of them done with Adaptive Optics in the near-IR) have started to find evidence for radial outflows in the NLR of Seyfert galaxies (Garc\\'ia-Lorenzo et al. 2001; Barbosa et al. 2009; Storchi-Bergmann et al. 2010; Riffel et al. 2010).  %suggests that the NLR gas accelerates outward from the nucleus to a turnover %point, and subsequently appears to decelerate, there exist %to the systemic velocity.  %The luminous central engine in active galaxies, most likely an %accreting supermassive black hole (BH), ionises the surrounding %interstellar medium leading to the broad-line region (BLR) %in the inner part and the narrow-line region (NLR) further out. %Among the central AGN components, the NLR has the advantage %to be directly accessible via imaging and spatially resolved %spectroscopy, at least for nearby AGNs.  %The narrow-line region (NLR) of active galactic nuclei (AGNs) %is an important laboratory to probe the interaction between the %nuclear engine and the circumnuclear interstellar medium (ISM) %of the host galaxies.  %With the launch of the Hubble Space Telescope (HST ) with its high angular %resolution (0B1), the NLR of Seyfert galaxies has received considerable attention. %With the long-slit capability of the Space Telescope Imaging Spectrograph (STIS) %and the Faint Object Camera (FOC), detailed constraints on kinematic models of the %NLR in the two types of Seyfert galaxies became possible. In turn, these models %provide good diagnostics on which dynamical analyses can be based. Biconical %outflow modeling of the NLR using STIS data was previously done by our group for %NGC 1068 (Seyfert 2), Mrk 3 (Seyfert 2), and NGC 4151 (Seyfert 1).  The high-ionization lines, or coronal lines, observed in the spectra of many AGNs, are forbidden fine structure transitions in the ground level of highly ionized atoms which have threshold energies above 100 eV. %The study of coronal lines is crucial because they %Accordingly, the coronal lines are unique tracers of AGN activity. In contrast to [O~{\\sc iii}] or H$\\alpha$, which are commonly used to study the NLR, the coronal lines are free of potential contributions from star formation. %shock-ionized gas, and tidal tails. Therefore, these lines provide direct insights into the structures and dynamical forces associated with the AGN. Ground-based spectroscopic studies revealed that these lines are present with comparable strength in AGN spectra regardless of their class (Penston et al. 1984; Erkens et al. 1997; Prieto \\& Viegas 2000; Rodr\\'iguez-Ardila et al. 2002; Reunanen et al. 2003), tend to be broader than NLR lines ($400<$FWHM$<1000$ km s$^{-1}$), and their centroid velocity is often blueshifted with respect to the systemic velocity of the galaxy (e.g. Rodr\\'iguez-Ardila et al. 2006). These facts have been interpreted as evidence for a location of the coronal-line region (CLR) closer to the AGN than the NLR, but outside the broad line region (BLR), and probably associated with outflows (e.g. M\\\"uller-S\\'anchez et al. 2006, hereafter MS06). Because of their high ionization potential (IP$>100$ eV), %associated with CLR emission, highly energetic processes are required. the lines can be produced by either a hard UV to soft X-ray continuum, fast shocks, or a combination of both. The success of photoionisation models to reproduce, within a factor of few, the measured line ratios points towards photoionization as the main excitation mechanism of the CLR \\citep{kraemer00, groves04, mullaney08, mazzalay10}, although an additional contribution from fast shocks must be invoked to fully explain the line ratios in several cases (MS06; Rodr\\'iguez-Ardila et al. 2006; Contini et al. 2009). However, the morphology, size and kinematics of the CLR remain uncertain because in most observational studies the coronal lines are not spatially resolved.  Detailed studies of the physical conditions across the CLR were not possible before the advent of the $HST$ and ground-based Adaptive Optics (AO). Recently, \\citet{prieto05} determined, for the first time, the size and morphology of the CLR in four Seyfert galaxies. These authors found that the CLR as traced by the [Si~{\\sc vii}] $\\lambda 2.48 \\micron$ line extends out to $r=30\u0096-200$ pc, and is aligned preferentially with the direction of the lower ionization cones seen in [O~{\\sc iii}] or H$\\alpha$. %As expected, the CLR is primarily produced very close %to the active nucleus but is nevertheless resolved over several %tens of parsecs. The morphology and size of the CLR indicate that the nuclear radiation is either intrinsically collimated or that some additional excitation mechanism (such as fast shocks) contributes to the formation of coronal gas at those observed distances from the ionizing source. With the long-slit capability of $HST$ (STIS), \\citet{mazzalay10} attempted to study the CLR kinematics in a sample of ten Seyfert galaxies. They found that the radial velocity curves of the coronal lines follow the same pattern as those of  [O~{\\sc iii}] emission and concluded that neither the strength nor the kinematics of the CLR scale in any obvious and strong way with the radio jet.  %of NGC1068 and concluded that the radio jet is not the principal %driving force on the outflowing NLR clouds.  %Due to the fact that the CLR is more compact than the classical NLR. %Nevertheless, the angular sizes of the CLRs in Seyfert galaxies are quite small.  %They are observed in the spectra of many AGNs, and appear to be approximately %equally abundant in type 1 and type 2 AGN (Rodriguez-Ardila et al., in prep).  %The study of coronal lines is crucial because they are unique tracers of AGN %activity. In contrast to [O III] or H$\\alpha$ which are commonly used to study the %NLR, the coronal lines are free of potential contributions from star formation, %shock-ionized gas, and tidal tails. Therefore, the CLR provides direct insights %into the structures and dynamical forces associated to the AGN. For example, the %CLR can provide clean and clear evidence that the bulk of the ionized gas close to %the AGN is in an outflow.  %So far, the nature, kinematics and origin of the CLR remain uncertain, as in most %observational studies the coronal lines are not spatially resolved. %If the excitation occurs via collisions, the gas %temperature should be about T = 106 K. In the case of photoionization via the hard %AGN continuum, temperatures of only a few 103 \u00a1 104 K are needed %Powerful pure starbursts, may show [He ii] %lines but not lines of higher IP ions (IP > 54 eV). %CLs have been detected in the optical and infrared spec- %tra of all types of Seyfert galaxies, including type 1s, type %2s, and narrow-line Seyfert 1 galaxies (e.g., Seyfert 1943; %Penston et al. 1984; Marconi et al. 1994; Nagao, Taniguchi %& Murayama 2000; Sturm et al. 2002; Rodr\u00b4\u00fdguez-Ardila et %al. 2002, 2006; Deo et al. 2007; Mullaney & Ward 2008; Ko- %mossa et al. 2008; Gelbord, Mullaney & Ward 2009) and %also radio galaxies (e.g., Best, Rottgering & Lehnert 1999; %Holt, Tadhunter & Morganti 2006), and represent one of the %key gaseous components of the active nucleus.  %Integral-field spectroscopy is an ideal tool for studying objects with complex %morphologies and kinematics, such as the NLR and CLR of Seyfert galaxies. %Recent studies based on ground-based integral-field spectroscopy (some of them %coupled to an AO module) have found evidence for radial outflows in the NLRs of %Seyfert galaxies (Barbosa et al. 2009; Riffel et al. 2009; Storchi-Bergmann et al. %2009).  There is great current interest in NLR and CLR kinematics, to resolve the much debated question of feedback from AGN via outflows of ionized gas, now recognized as a potentially crucial input to the $M_{BH}-\\sigma^*$ relation (e.g. Gebhardt et al. 2000): AGN outflows provide energy into the ISM of the host galaxy, controlling the star formation and accretion processes, thus regulating the black hole and bulge growth. Recent results from Krongold et al. (2007) and Storchi-Bergmann et al. (2010) indicate that the total kinetic energy released by the AGN outflow is comparable to the energy required to disrupt the hot phase of the ISM, which could then disrupt star formation. These important topics had not previously been widely studied because high signal-to-noise (S/N) spectra over a full two-dimensional (2D) field with high spatial resolution are required, and these are still scarce. Particularly, as the coronal lines are exclusively driven by the AGN power, the CLR is an ideal laboratory for investigating the physical processes within $<300$ pc from the central black hole.  %Seyfert galaxies are important laboratories to study in detail the co-evolution of %supermassive black holes (SMBHs) and their host galaxies. The reason is simply %that these are the nearest objects in which the NLRs and CLRs are directly %accessible via spatially resolved imaging and spectroscopy. In an effort to draw general conclusions about the kinematics of the NLR and CLR in active galaxies, %In order to draw general conclusions about the physical properties and kinematics %of the NLR and CLR in active galaxies, we have carried out a program of high spatial resolution, near-IR integral field spectroscopy of Seyfert galaxies. The measurements discussed here were obtained with the use of AO and the instruments SINFONI on the European Southern Observatory (ESO) 8.2 m Very Large Telescope (VLT) and OSIRIS on the 10 m Keck II Telescope. In this paper, which is the first of a series devoted to investigate the role of NLR and CLR kinematics in the co-evolution of supermassive black holes (SMBHs) and their host galaxies, %the application of ionized gas 2D kinematics to study in detail the co-evolution %of SMBHs and their host galaxies, we discuss the NLR and CLR distribution and kinematics of seven Seyfert galaxies using the Br$\\gamma$ $\\lambda 2.16 \\micron$ recombination line as tracer of the NLR and the near-IR coronal lines:  [Si~{\\sc vi}] $\\lambda 1.96\\micron$ (IP $= 167$ eV),  [Al~{\\sc ix}] $\\lambda 2.04\\micron$ (IP $= 285$ eV) and [Ca~{\\sc viii}] $\\lambda 2.32\\micron$ (IP $= 127$ eV). The existing studies of the kinematics of the NLR are mostly based on long-slit spectroscopy obtained in the optical with $HST$. Our goal is to extend these studies to the near-IR, a region which is less affected by dust, and where AO permits diffraction-limited integral field spectroscopy of various emission lines that are a vital addition to the lines observed at optical and UV frequencies. Furthermore, this work serves as the basis for further studies of the NLR with future near-IR telescopes and interferometers, such as, e.g. LBT, ELT and $JWST$. We aim to determine for the first time the 2D kinematics of the CLR, examine previous claims of radial outflow along the NLR, as well as investigate possible interactions between the NLR/CLR and other components of the AGN or the host galaxy in the central $<300$pc. When a significant outflowing structure is present, the 2D kinematics are characterized with appropriate kinematic models and the mass outflow rates are calculated. %is calculated and its corresponding kinetic power is compared with the bolometric %luminosity of the AGN to gauge the impact of the outflow.  %Spatially resolved integral-field observations of the Br$\\gamma$ $\\lambda 2.17 %\\micron$m recombination line and the near-IR coronal lines: [SiVI] $1.96\\micron$m, %[AlIX] $1.96\\micron$m  and [CaVIII] $\\lambda 2.32\\micron$m will show the extent %and kinematics of the NLRs and CLRs. The paper is structured as follows. In Section 2 we describe the sample characteristics, observations and the methods used for the data reduction and analysis. The general properties of the NLR and CLR are described in Section 3. In Section 4 we analyze in detail the observed kinematics and perform kinematic modeling. The implications of the results are discussed in Section 5, including the origin of the ionized gas motions and possible interactions with other components of the unified model of AGN and the host galaxy. The overall conclusions of the study are outlined in Section 6, and in the Appendix a more detailed description of the individual objects is presented. Throughout this paper, $H_0 = 75$ km s$^{-1}$ Mpc$^{-1}$ is assumed.   %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. ",
        "conclusions": "\\label{conclusions}  SINFONI and OSIRIS near-IR AO-assisted integral field spectroscopy has been used in this work to probe the kinematics of the NLR and CLR in nearby AGNs. %most of which are Seyfert 1 galaxies. Thanks to our high S/N integral field data we were able to extract the 2D distribution and kinematics of Br$\\gamma$ and the high-ionization line [Si~{\\sc vi}] in the central $<300$ pc with spatial resolutions ranging from 4 to 36 pc. Our main results and conclusions can be summarized as follows:  \\begin{itemize} \\item In all cases the NLR and CLR are resolved (except in NGC~3783), and their morphologies are similar, indicating that they are being produced by the same ionizing source. The two regions present a bright compact core and extended emission. While the size of the NLR seems to be larger than our FOV in some objects, the CLR is more compact and could be well measured in the sample galaxies (except in Circinus) with sizes ranging from $r=80$ pc in NGC~6814 to $r=150$ pc in NGC~1068. \\item The kinematics of the NLR at these scales can be put into three groups: (a) velocity fields dominated by rotation, (b) disturbed rotational patterns, and (c) velocity fields dominated by non-circular motions. For the galaxies with velocity fields dominated by rotation, the NLR and stellar velocity fields are similar and indicate that the majority of the NLR gas is distributed in a disk. \\item The velocity and dispersion maps of the NLR and CLR in the sample galaxies are similar, with a trend for the low-ionization gas to be a better tracer of the disk component, while the high-ionization gas presents more deviant behavior with velocity fields dominated by non-circular motions. \\item Several pieces of evidence suggest that the non-circular motions observed in the velocity fields of the NLR and CLR correspond to outflows. The main argument favouring this interpretation is the evidence for radial acceleration to a projected distance of $\\sim100$ pc (followed by deceleration in some cases) observed in the velocity channel maps and the radial velocity maps. Other evidences are: (1) For the galaxies with existing radio images, the velocity channel maps reveal blueshifts in the direction of the radio jet pointing towards us and redshifts in the direction of the jet pointing away from us, (2) the velocity dispersion clearly increases at the nucleus and/or along the directions of the blueshifts and redshifts and (3) the non-circular velocities are too high to be explained by any sort of reasonable gravitational potential. \\item Kinematic models consisting of rotation and/or biconical outflow have been successfully applied to the data. The models provide a good match to the overall velocity patterns, despite slight local variations in the velocity fields. \\item Regarding the rotational component, the data clearly shows that the disks of low- and high-ionized gas are spatially coincident with the inner disks of molecular gas, suggesting that the three types of gas are spatially mixed. Since high excitation emission is observed outside the borders of the bicone, the torus must be a clumpy medium such that a fraction of ionizing photons escape in different directions and ionize the gas in the disk. %Furthermore, since the ionized gas emission extends beyond the borders of the %bicones as represented by the models, the molecular gas within this nuclear region %is in a clumpy distribution such that a fraction of ionizing photons escape in %different random directions outside the bicone and ionize the gas in the disk. %Furthermore, since the high excitation emission is observed outside the borders of %the bicone, the molecular gas whithin %the torus must be a clumpy medium in order to allow the escape of radiation along %the equatorial plane of the bicone and ionize the gas in the disk. %Furthermore, the ionized gas emission is observed also along its equatorial plane, %indicating that the AGN ionizes gas beyond the limits of the bicone. \\item The maximum velocity of the outflowing component ranges between 120 km s$^{-1}$ in NGC~6814 to 1900 km s$^{-1}$ in NGC~1068, and it is reached at a typical distance of $\\sim180$ pc from the AGN. We observe two types of outflows based on their opening angle: collimated with full opening angles $<70\\degr$ and less collimated with full opening angles $>90\\degr$. \\item We find an anticorrelation between the outflow velocity ($v_{\\mathrm{max}}$) and the half-opening angle of the bicone ($\\theta_{\\mathrm{out}}$), and another anticorrelation between $\\theta_{\\mathrm{out}}$ and the molecular gas mass ($M_{\\mathrm{gas}}$) in $r<30$ pc. These anticorrelations suggest that galaxies with greater amount of gas in their centers have higher degree of collimation (small $\\theta_{\\mathrm{out}}$) and higher outflow velocities. This indicates that the torus invoked by unified models is indeed the structure responsible for the collimation of the outflows on small spatial scales and that the accumulation of gas clouds towards the nucleus (higher $M_{\\mathrm{gas}}$) increments its collimation power and produces outflows with higher velocities (stops the outflow expanding sideways). %In addition, it must be a clumpy medium in order to allow the escape of radiation %along the equatorial plane of the bicone and ionize the gas in the disk. %producing the observed flux and velocity maps. \\item\tThe observed bipolar outflows are consistent with the torus model of AGN (except in NGC~3783). %if we consider a bicone with internal and external faces, i.e. it has a thickness. The five intermediate type Seyfert galaxies studied here (four Sy 1.5s and one Sy 1.9) are dominated by blueshifts on one side and redshifts on the other side of the AGN, consistent with an intermediate viewing angle, whereas the Seyfert 2 galaxy NGC 1068 shows blueshifts and redshifts on either side of the AGN, consistent with the unified model prediction of a bicone axis near the plane of the sky. Considering the uncertainties in the kinematic modeling, and possible misalignments in the torus-bicone system, it is likely that we have a direct view to the BLR in most of the intermediate type Seyferts (except in NGC~3783). We conclude that the majority of these objects (and possibly most Seyfert 1 galaxies) are actually seen close to the edge of the cone, rather than down the middle. Alternatively, a clumpy torus may explain the presence of broad lines even if the observer is not looking into the bicone, which is likely the case in NGC~3783. \\item We have estimated mass outflow rates in the range 1 to 10 M$_\\odot$ yr$^{-1}$, which are $\\sim10^2-10^3$ times the accretion rate necessary to feed the AGN. The estimated kinetic powers of the outflows are $10^2-10^4$ times smaller than the AGN bolometric luminosities. We note that $\\dot{E}_{\\mathrm{out}}$ in half of the galaxies presenting outflows (NGC~1068, NGC~2992 and NGC~4151) fulfill the requirements of two-stage feedback models to be thermally coupled to the circumnuclear gas ($\\dot{E}_{\\mathrm{out}}\\approx 0.005L_{\\mathrm{bol}}$), and match the $M_{BH}-\\sigma^*$ relation. %If the density of the CLR has been underestimated, it is plausible that %$\\dot{M}_{\\mathrm{out}}$ could be at least one order of magnitude higher, %resulting in $\\dot{E}_{\\mathrm{out}}\\approx 0.005L_{\\mathrm{bol}}$ in all %galaxies. The two galaxies with the highest $\\dot{E}_{\\mathrm{out}}/L_{\\mathrm{bol}}$ (NGC 1068 and NGC 4151) are also the only objects in which a clear interaction of the surrounding medium with the radio jet is observed. This may suggest that Seyfert galaxies with strong and collimated radio emission are also hosts of powerful outflows of ionized gas. \\item The high velocity components detected at the nucleus of some galaxies and the high $\\dot{M}_{\\mathrm{out}}/\\dot{M}_{\\mathrm{acc}}$ ratios indicate that the outflowing gas is not likely a single wind from the AGN, but the result of a two-stage process in which the circumnuclear ISM has been pushed away by an already accelerated AGN wind. The radio jet may also have a contribution to the formation of the outflows at small spatial scales, transfering momentum to the circumnuclear gas. %Our observations discard the radio jet as the main driving mechanism of the large %scale NLR/CLR outflows, as well as the post-starburst detected in the nuclear %regions of these galaxies. We propose that the observed bipolar outflows are the %result of a two-stage process involving the local interaction of high velocity gas %(an accelerated wind from the AGN) with its ambient medium. \\end{itemize}  %This study demonstrated the power of integral field spectroscopy for the determination of kinematic centers of galaxies.  %Further AO-assisted near-IR integral-field observations of Seyfert 2 galaxies %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"
    },
    "1107/1107.0591_arXiv.txt": {
        "abstract": " ",
        "introduction": "} The study of the properties of young massive star clusters in our Local Group can be done by means of the census of the stars that belong to the corresponding object and the analysis of their properties using a colour-magnitude diagram. In the case of those objects whose stellar population cannot be resolved with the facilities that are nowadays available, it is necessary to appeal to other tecniques. A possible approach is to study the massive stars by inspecting their effects on the surrounding interstellar medium (ISM), including the supply of mechanical energy from stellar winds and supernovae explosions, the ejection of new metals produced in the interiors of the stars, or the excitation and ionization of the atoms in the ISM by the energetic photons emitted by the stars. These factors make the optical spectrum of an H{\\sc ii} region to be frequently characterized by the presence of bright prominent recombination lines emitted by hydrogen and helium and of collisional lines emitted by the metals. Although the shape of the stellar continuum and some stellar features are sometimes detectable ({\\em e.g.} the Wolf-Rayet features), the most important source of information about the properties of the ionizing stellar clusters can only be the relative intensities of these lines.  These intensities are mainly dominated by the so-called functional parameters, including the metallicity ($Z$), the ionization parameter ($U$, {\\em i.e.} the ratio between ionizing photons and the density of particles), and the equivalent effective temperature ($T_*$). This last is the unique parameter that only depends on the spectral energy distribution (SED) of the ionizing star or cluster. ",
        "conclusions": "The determination of properties of a massive young stellar cluster can be obtained by using information from the emission line spectrum from the ionized gas in those cases where the stellar population of the cluster cannot be resolved. The $\\eta$ parameter as analyzed in 2 dimensions is a robust method to derive the equivalent effective temperature of a cluster, even in those cases where there are variations in the geometry of the gas or in the dust-to-gas ratio. The most important caveat appears as a consequence of the disagreements in the involved energetic ranges between different synthesis model atmospheres, but it is possible to use this to study the relative variations of $T_*$ in homogeneous samples of objects. This is the case of the variation of $T_*$ across spiral disks found by P\\'erez-Montero \\& V\\'\\i lchez (2009), which points to a hardening of the ionizing radiation above all in spiral galaxies of low dynamical mass and luminosity with late morphological type.    \\small  %"
    },
    "1107/1107.4050_arXiv.txt": {
        "abstract": "In 2004, \\xsrc underwent a period of intense and long-lasting burst activity that included the giant flare of 27 December 2004 -- the most intense extra-solar transient event ever detected at Earth. During this active episode, we routinely monitored the source with Rossi X-ray Timing Explorer and occasionally with Chandra. During the course of these observations, we identified two relatively bright bursts observed with Konus$-$Wind in hard X-rays that were followed by extended X-ray tails or afterglows lasting hundreds to thousands of seconds. Here, we present detailed spectral and temporal analysis of these events observed about 6 and 1.5 months prior to the 27 December 2004 Giant Flare. We find that both X-ray tails are consistent with a cooling blackbody of constant radius. These spectral results are qualitatively similar to those of the burst afterglows recorded from SGR 1900+14 and recently from SGR 1550$-$5418. However, the latter two sources exhibit significant increase in their pulsed X-ray intensity following the burst, while we did not detect any significant changes in the RMS pulsed amplitude during the \\xsrc events. Moreover, we find that the fraction of energy partitioned to the burst (prompt energy release) and the tail (afterglow) differs by an order of magnitude between SGR 1900+14 and \\xsrcnos. We suggest that such differences can be attributed to differences in the crustal heating mechanism of these neutron stars combined with the geometry of the emitting areas. ",
        "introduction": "Soft Gamma Repeaters (SGRs) are a small group of isolated neutron stars that exhibit extraordinary properties. They are persistent X-ray sources with luminosities ranging between 10$^{33}$ and 10$^{35}$ erg s$^{-1}$, spin periods $2-9$ s and large spin down rates indicating extremely high surface dipolar magnetic field strengths of B$\\sim$10$^{14}$$-$10$^{15}$ G. They are distinguished from the overall neutron star population by their repeated emission of short and intense hard X-ray / soft gamma ray bursts with a typical duration of $\\sim0.1$ s. The burst repetition timescale is of the order of a few seconds with burst peak luminosities, $L_{\\rm peak}$, ranging between $10^{36}$-$10^{41}$ erg s$^{-1}$ \\citep{gogus99, gogus00}. The burst fluence distribution is continuous and can be well fitted with a power law of index of around -1.6 \\citep{gogus99,gogus00}. Consequently, although the short and less energetic events constitute the bulk of the SGR burst population, several intermediate size events (durations up to tens of seconds, $L_{\\rm peak} \\sim 10^{42}$ erg s$^{-1}$) have also been recorded in the last twelve years \\citep{mazets99,kouv01,feroci03,mereg09}. Finally, SGRs emit giant flares, which are exceptionally rare bright events ($10^{44} < L_{\\rm peak} <$ few$\\times 10^{47}$ erg s$^{-1}$) that last several hundreds of seconds. Only three such events have been recorded to date; their durations, temporal structure and energetics are very similar (for a detailed review, see Woods \\& Thompson 2006 or Mereghetti 2008).  An emerging class of SGR bursts referred to as intermediate bursts are always associated with afterglows (tails) when observed with sensitive, low-background X-ray detectors.  Lenters et al. (2003) compared three such bursts from \\sgrc along with the giant flare of 27 August 1998, and estimated that in all four cases the energy emitted during the afterglow (2$-$10 keV) was $\\sim 2$\\% of the total energy of the burst (25$-$100 keV). Their detailed spectral analysis of the tail portion of two of these events revealed a decaying thermal component, possibly associating the tails with the cooling of a burst-induced heating process on the surface of the neutron star \\citep{ibra01,lent03}. Unfortunately, there was no sensitive soft X-ray coverage of either the 1979 March 5 Giant Flare from SGR $0526-66$ or the June 1998 intermediate bursts from SGR $1627-41$. Recently, intermediate bursts from SGR 1550$-$5418 (aka 1E 1547.0$-$5408) were observed with multiple instruments during its 2009 burst active episode \\citep{mereg09, kaneko10, eneto10}.  One of the most prolific sources, \\xsrc, entered an active phase in March 2004 during which it emitted more than a thousand bursts over the following several months including the energetic Giant Flare of 2004 December 27 which had a peak luminosity of a few$\\times$$10^{47}$ erg s$^{-1}$  \\citep{hurl05,fred07}. We searched our {\\it Rossi X-ray Timing  Explorer} (RXTE) observations of the source between January 2004 $-$ May 2005  and identified only two intermediate type events out of a total of over 1500 bursts. Each was followed by an extended X-ray tail or afterglow lasting several hundred seconds.  We were very fortunate to obtain simultaneous coverage of one of these bursts with {\\it Chandra X-ray Observatory} (CXO). We  present here a detailed temporal and spectral analysis of the main burst and the afterglows of these two intermediate events (Sections 2 and 3), discuss their similarities and differences with other magnetar afterglows, and the implications thereof in Section 4. ",
        "conclusions": "We have identified two \\xsrc bursts during the source's extremely active episode in 2004, that exhibited extended tail emissions similar to the ones reported from SGR~$1900+14$ \\citep{lent03}. These events happened roughly at 6 and 1.5 months prior to the SGR Giant Flare on 2004 December 27. Earlier, G\\\"o\\u{g}\\\"u\\c{s} et al. (2000) identified 290 short bursts in the RXTE observations of \\xsrc during its 1996 active episode. Several of these events were longer than average (1$-$2 s) but none displayed any extended tail emission.  There are characteristic differences between two bursts we investigated here: the 2004 October 17 bursts is among the events with the highest energy fluence detected during the 2004 active episode of \\xsrc and shows clear spectral evolution. The fluence of the 2004 June 22 burst, on the other hand, was not as high and did not exhibit any spectral variation during the course of the event. Due to the fact that both events displayed extended tails following them, the condition for an extended X$-$ray afterglow emission to follow cannot be singly attributed to the energetics of the burst.  SGR~$1900+14$ emitted two strong bursts on 1998 August 29 and 2001 April 28 with extended tail emission (durations $\\sim$8000 s and $\\sim$5000 s, respectively). Lenters et al. (2003) studied these event tails extensively; they report that the burst fluences (25$-$100 keV) are 1.9$\\times$10$^{-5}$ erg cm$^{-2}$ and 8.7$\\times$10$^{-6}$ erg cm$^{-2}$, respectively. To directly compare these to the \\xsrc events, we used the Konus-Wind spectral results of the main pulse to estimate the fluences of the latter between $25-100$ keV; these are 3.6$\\times$10$^{-6}$ erg cm$^{-2}$ and 4.3$\\times$10$^{-5}$ erg cm$^{-2}$, for the events of 2004 June 22 and October 17, respectively. Therefore, the main burst energetics of these intermediate events from both SGRs are quite similar in the energy range where most of their photons are emitted ($25-100$ keV).  Lenters et al. (2003) also found that there was a significant correlation between the SGR~$1900+14$ event tail fluences in the 2$-$10 keV range and the leading main burst fluences in 25$-$100 keV. They estimated the former to be about 2\\% of the latter fluence.  To check whether a similar correlation is present in the \\xsrc events, we determined the ratios of their tail fluences in the 2$-$10 keV range (1.2$\\times$10$^{-8}$ erg cm$^{-2}$ and 2.7$\\times$10$^{-7}$ erg cm$^{-2}$) to these in their main events. These are 0.34\\% and 0.63\\% for the June 22 and October 17 events, respectively. Although we examined here only two events, we note that these ratios for the \\xsrc bursts are both of the same order of magnitude and 3$-$6 times smaller than those of the SGR~$1900+14$ bursts.  The tail spectra of the two events from SGR~$1900+14$ also included thermal components; the dimmer one (2001 April 28 event) was exclusively thermal and consistent with a BB temperature of 2.35$\\pm$0.05 keV, while the brighter one (1998 August 29) had both a thermal (kT = 2.2$\\pm$0.2 keV) and a non-thermal (power law index of 2.1$\\pm$0.2) component (Lenters et al. 2003). Similarly, we find that the dimmer \\xsrc event tail can be described by a single thermal component with a temperature of 3.4$\\pm$0.2 keV, while the brighter one required two spectral components (kT = 3.5$\\pm$0.3 keV and photon index = 1.6$\\pm$0.1). The BB temperatures obtained from the two \\xsrc events are consistent with each other, and significantly larger than those of the SGR~$1900+14$ events. The dependance of the power law component on the burst tail intensity can also be demonstrated from Figure \\ref{fig:tflux_tail2}: There we see that the power law component dominates at the later--dimmer stages of the tail, which are only visible in brighter events. As the events become dimmer themselves, only the BB spectral component can be seen in the data.  \\citet{espo07} performed a time resolved spectral analysis of the first 7.5 hours of the afterglow emission following the 2001 April 18 flare from SGR 1900+14. They found that an additional BB component was required to significantly better fit the afterglow spectrum. They obtain the best fit parameters by keeping the radius of the BB emitting region constant at 1.6 km and allowing the BB temperature vary. The temperature of the additional BB component gradually declines from 1.23 to 0.92 keV over the course of 7.5 hours. They concluded similarly that the cooling thermal emission from a constant region could account for the additional spectral component.  Recently, \\citet{kaneko10} identified a unique enhancement of hard X-ray emission following an intense burst from SGR J1550$-$5418 at the onset of its major outburst episode in 2009. The enhanced emission showed clear pulsations with a period consistent with the spin period of the source. Detailed temporal study revealed that the pulsation was strongest around the peak of the enhancement time-wise and in 50--74\\,keV energy-wise with a pulsed fraction of $\\sim$55\\%. Hard X-ray spectra of the enhancement were well described with a two-component model: a power law with a BB. While the temperature of the BB ($\\sim$17\\,keV) remained constant throughout the enhancement period, the BB component was most evident around the peak time, where the pulsation was also strongest. The estimated total energy emitted during the enhancement (i.e., the afterglow) was $2.9 \\times 10^{40}$\\,erg, of which the BB accounted for 19\\%. Based on the BB flux, the emitting radius of the thermal component could be as small as 0.12 km. One possible interpretation is that a small region on the surface of the neutron star is likely heated by hot plasma confined by the twisted magnetic field, in which case the energy can be dissipated as the magnetic field untwists \\citep{belo09}.  The pulsed amplitudes of the persistent X-ray emission from SGR 1900+14 during the extended tails were remarkably larger than before the bursts. In fact, the periodic oscillations from SGR 1900+14 were clearly visible in the decaying tail of both intermediate events \\citep{ibra01,lent03}, as well as during the enhanced hard X-ray emission of SGR J1550$-$5418 \\citep{kaneko10}. In the case of \\xsrc events, the changes in the pulsed amplitude in the tails of intermediate bursts were not significant. Moreover, the fraction of energy released in the tail (afterglow) relative to that released during the main the burst (prompt energy release) differs by about an order of magnitude between SGR 1900+14 and \\xsrcnos. Although the spectral properties of the \\xsrc extended tails are similar to those of the burst afterglows recorded from SGR 1900+14 and from SGR J1550$-$5418, the constancy of the pulsed amplitude in \\xsrc and the difference in the fractional energy release indicate intrinsic differences in burst-induced crystal heating (and cooling) mechanisms in these three sources.  We also note here another significant difference between the persistent emission behaviors following Giant Flares from two different sources. The energy released during the 2004 December 27 Giant Flare from \\xsrc was at least two orders of magnitude larger than the 1998 August 27 Giant Flare from SGR 1900+14. However, the persistent X-ray flux of \\xsrc resumed quickly toward its long term value in the aftermath of the flare \\citep{rea05,woods07}, while the Giant Flare of SGR 1900+14 was followed by a long lasting ($\\gtrsim$100 days) flux enhancement \\citep{woods01}. In the light of these facts, we suggest that \\xsrc can efficiently cool by radiation (and baryonic material release \\citet{gaensler05}) during the burst (prompt release); most likely also comparatively less energy is imparted to the deep crustal heating than was the case for SGR$1900+14$. An additional possibility for the lack of significant pulsed intensity increase in \\xsrc may be the geometry of the emitting areas: the heated crustal surface might be close to the rotation axis, which could be nearly aligned to the magnetic axis of the neutron star. This possibility is indirectly supported by the fact that the main peak of the 2004 October 17 intermediate burst encompasses the peak of its pulse profile."
    },
    "1107/1107.5056_arXiv.txt": {
        "abstract": "We consider the interaction between a binary system (e.g. two supermassive black holes or two stars) and an external accretion disc with misaligned angular momentum. This situation occurs in galaxy merger events involving supermassive black holes, and in the formation of stellar--mass binaries in star clusters. We work out the gravitational torque between the binary and disc, and show that their angular momenta $\\mathbi{J}_{\\rm{b}}, \\mathbi{J}_{\\rm{d}}$ stably counteralign if their initial orientation is sufficiently retrograde, specifically if the angle $\\theta$ between them obeys $\\cos\\theta < -J_{\\rm{d}}/2J_{\\rm{b}}$, on a time short compared with the mass gain time of the central accretor(s). The magnitude $J_{\\rm{b}}$ remains unchanged in this process. Counteralignment can promote the rapid merger of supermassive black hole binaries, and possibly the formation of coplanar but retrograde planets around stars in binary systems. ",
        "introduction": "\\label{intro}  Galaxy mergers are commonly thought to be the main mechanism driving the coevolution of galaxies and their central supermassive black holes (SMBH). In such a merger we expect the formation of a SMBH binary in the centre of the merged galaxy. Gravitational waves quickly drive the binary to coalesce if the orbital separation can be shortened to $\\lesssim 10^{-2}$pc. The binary may stall at a separation greater than this if the interaction with the merged galaxy is not efficient enough in extracting orbital angular momentum and energy. For the stellar component of the galaxies this occurs at approximately a parsec, creating ``the final parsec problem'' \\citep{MM2001}. There have been many papers exploring potential solutions to this problem, for example a sling-shot mechanism involving a triple SMBH system \\citep{Iwasawaetal2006}, efficient refilling of the binary loss cone by angular momentum exchange between stellar orbits and a triaxial dark matter halo \\citep{Bercziketal2006} and also the evolution of the binary with a prograde accretion disc (circumbinary discs: \\citealt{AN2005}; \\citealt{MM2008}; \\citealt{Lodatoetal2009}; \\citealt{Cuadraetal2009} and embedded discs: \\citealt{Escalaetal2005}; \\citealt{Dottietal2007,Dottietal2009}). In a recent paper \\citep{Nixonetal2011} we explored the evolution of a binary interacting with a {\\it retrograde} circumbinary accretion disc. We showed that this is more efficient than a prograde disc in removing binary orbital angular momentum and energy. This is simply because there are no orbital resonances between the binary and the disc and thus there is direct accretion of retrograde gas onto the binary.  Here we consider the alignment process between a binary system and an external misaligned accretion disc. This situation can arise in at least two astronomical contexts. First, a merger event between galaxies can produce a SMBH binary in the centre of the merged galaxy, and this or a later accretion event may surround the hole with a disc of accreting gas. A similar situation arises during the formation of stars in a cluster. A binary system may form, but also capture gas into an external disc.  In both of these cases, there is no compelling reason to assume that the binary and disc rotation are initially parallel or even roughly coaligned (cf \\citealt{KP2006}). As we shall see, the gravitational interaction between the binary and the disc generates differential precession in the disc gas, and thus viscous dissipation. This gives a dissipative torque which vanishes only when the binary and disc angular momenta $\\mathbi{J}_{\\rm{b}}, \\mathbi{J}_{\\rm{d}}$ are either parallel or antiparallel. In all such cases, the torque diffuses the tilt or warp through the disc (cf \\citealt{Pringle1992,Pringle1999}; \\citealt{WP1999}) driving the system to one of these equilibria. The existence of a warp makes the precise definition of disc angular momentum $\\mathbi{J}_{\\rm{d}}$ quite subtle and we return to this point in the Discussion.  The binary--external disc interaction is very similar to the effect of the Lense--Thirring (LT) precession on an accretion disc around a spinning black hole (\\citealt{BP1975}; \\citealt{Pringle1992}; \\citealt{SF1996}; \\citealt{NP1998}; \\citealt{AN1999}; \\citealt{NA1999}; \\citealt{NP2000}; \\citealt{LP2006} etc) if we replace $\\mathbi{J}_{\\rm{b}}$ by the hole spin angular momentum $\\mathbi{J}_{\\rm{h}}$. For some years it was thought that the LT interaction always led to co--alignment (i.e. $\\mathbi{J}_{\\rm{b}}$ and $\\mathbi{J}_{\\rm{d}}$ parallel). However \\citet{Kingetal2005} (hereafter KLOP) showed on very general grounds that counteralignment does occur, if (and only if) the initial angle $\\theta$ between $\\mathbi{J}_{\\rm{d}}$ and $\\mathbi{J}_{\\rm{h}}$ satisfies $\\cos\\theta < -J_{\\rm{d}}/2J_{\\rm{h}}$, where $J_{\\rm{d}} = \\left|\\mathbi{J}_{\\rm{d}} \\right|$ and $J_{\\rm{h}} = \\left|\\mathbi{J}_{\\rm{h}} \\right|$. \\citet{SF1996} had implicitly assumed $J_{\\rm{d}} \\gg J_{\\rm{h}}$ and so enforced co--alignment. With this restriction lifted, \\citet{KP2006,KP2007} and \\citet{Kingetal2008} showed that accretion from a succession of randomly--oriented discs leads to spindown of the supermassive black hole, allowing rapid mass growth.  In this paper we examine the alignment process for a binary and an external disc. We show that the argument of KLOP is generic, and that the disc and binary counteralign if and only if $\\cos\\theta < -J_{\\rm{d}}/2J_{\\rm{b}}$. As a result it is quite possible for SMBH binaries to be surrounded by a completely retrograde disc which strongly promotes coalescence (cf \\citealt{Nixonetal2011}). In the case of a newly--formed stellar binary, the presence of a counteraligned disc can lead to the formation of planets with retrograde orbits. ",
        "conclusions": "\\label{discussion}  So far in this paper we have avoided fully spelling out the meaning of the disc angular momentum $\\mathbi{J}_{\\rm{d}}$. This is complicated because the binary torque falls off very strongly with radius, and so a large contribution to the angular momentum in a distant part of the disc may be irrelevant to the alignment process, or affect this process in a time--dependent way (cf \\citealt{LP2006}). Sections 3 and 4 of KLOP discuss these questions in more detail. Effectively $\\mathbi{J}_{\\rm d}$ can be thought of as the disc angular momentum inside the warp radius, and therefore a time--dependent quantity.  At early times $J_{\\rm{d}}$ is small, as only a fraction of the total gas interacts with the binary. Counter--alignment may occur if $\\theta > \\pi/2$, but at later times, as $J_{\\rm{d}}$ grows and more gas is able to interact with the binary, alignment eventually happens (when $J_{\\rm{d}} > 2J_{\\rm{b}}$). So if $\\theta >\\pi/2$, even for $J_{\\rm{d}} > 2J_{\\rm{b}}$ we expect $\\sim 2J_{\\rm{b}}|\\cos\\theta|$ of disc angular momentum to counteralign with the binary before the outer disc comes to dominate and enforce coalignment (cf. \\citealt{LP2006}).   The typical timescale for co-- or counter--alignment for a SMBH binary is \\begin{equation} t_{\\rm binary} \\simeq {J_{\\rm{b}}\\over J_{\\rm{d}}(R_{\\rm{w}})}{R_{\\rm{w}}^2\\over \\nu_2} \\label{talign} \\end{equation} where $R_{\\rm{w}}$ is the warp radius, $J_{\\rm{d}}(R_{\\rm{w}})$ is the disc angular momentum within $R_{\\rm{w}}$, and $\\nu_2$ is the vertical disc viscosity. This is identical to the formal expression for LT alignment of a spinning black hole if we replace the spin angular momentum $J_{\\rm{h}}$ with $J_{\\rm{b}}$ (cf \\citealt{SF1996}). The warp radius is given by equating the precession time $1/\\Omega_{\\rm{p}}(R)$ to the vertical viscous time $R^2/\\nu_2$ . Inside this radius the precession timescale is short and the disc dissipates and co-- or counter--aligns with the binary plane. Outside this radius the disc is not dominated by the precession and so maintains its misaligned plane. The connecting region therefore takes on a warped shape shown in \\fref{disc}. As time passes the warp propagates outwards and co-- or counter--aligns the entire disc with the binary plane.  \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=20pc]{disc_pic.eps} \\caption{The warped disc shape expected after the inner disc co-- or counter--aligns with the binary plane but the outer disc stays misaligned. Eventually the entire disc will co--or counter--align with the binary plane, depending on the global criterion (eqn~\\ref{cond}). Note that in practice precession makes the warp non--axisymmetric.} \\label{disc} \\end{center} \\end{figure}  Approximating the disc angular momentum as \\begin{equation} J_{\\rm{d}}(R_{\\rm{w}}) \\sim \\pi R_{\\rm{w}}^2\\Sigma(GMR_{\\rm{w}})^{1/2} \\label{jwarp} \\end{equation} with $\\Sigma$ the disc surface density and $M = M_1 + M_2$, and using the steady--state disc relation $\\dot M = 3\\pi\\nu\\Sigma$ we find \\begin{equation} t_{\\rm binary} \\sim 3{M_2\\over M_1}\\biggl({a\\over R_{\\rm{w}}}\\biggr)^{1/2}{\\nu_1\\over \\nu_2}{M\\over \\dot M}, \\label{talign2} \\end{equation} where we have also used \\begin{equation} J_{\\rm{b}} = M_1M_2\\biggl({Ga\\over M}\\biggr)^{1/2}. \\label{jb} \\end{equation} Since $\\nu_1 < \\nu_2$ \\citep{PP1983}, $a \\ll R_{\\rm{w}}$ and $M_2 < M_1$, we see that alignment takes place on a timescale shorter than the mass growth of the central accretor(s).  The timescale (\\ref{talign2}) is directly analogous to the expression \\begin{equation} t_{\\rm LT} \\sim 3a_*\\biggl({R_{\\rm{s}}\\over R_{\\rm{w}}}\\biggr)^{1/2}{\\nu_1\\over \\nu_2}{M\\over \\dot M} \\label{lt} \\end{equation} for alignment under the LT precession, where $a_* < 1$ is the Kerr spin parameter and $R_{\\rm{s}}$ the Schwarzschild radius of the spinning hole. Evaluating $R_{\\rm{w}}$ in the two cases we find \\begin{equation} {t_{\\rm LT}\\over t_{\\rm binary}} \\sim {3^{1/2}\\over 2}\\biggl({a_*\\over M_2/M_1}\\biggr)^{1/2}\\biggl({a\\over R_{\\rm{s}}}\\biggr)^{1/4}. \\end{equation} Thus in general, provided we assume that the ratio $\\nu_1/\\nu_2$ is similar in the two cases and that the hole spin is not rather small ($a_* < (R_{\\rm{s}}/a)^{1/2}(M_2/M_1)$), then the binary--disc alignment is rather faster than the corresponding process for spinning black holes.  Our result has significant consequences for SMBH binaries. For random orientations, \\eref{cond} shows that initial disc angles leading to alignment occur significantly more frequently than those giving counteralignment only if $J_{\\rm d} > 2J_{\\rm b}$. (In the LT case this fact leads to a slow spindown of the hole, because retrograde accretion has a larger effect on the spin, \\citealt{Kingetal2008}.) A number of studies (\\citealt{AN2005}; \\citealt{MM2008}; \\citealt{Cuadraetal2009}; \\citealt{Lodatoetal2009}) have shown that prograde external discs are rather inefficient in shrinking SMBH binaries and solving the last parsec problem. This is essentially because of resonances within the disc. In contrast, the slightly rarer retrograde events have a much stronger effect on the binary. These rapidly produce a counterrotating but coplanar accretion disc external to the binary, which has no resonances. We note that \\citet{Nixonetal2011} show that the binary gradually increases its eccentricity as it captures negative angular momentum from the disc, ultimately coalescing once this cancels its own. A non--zero binary eccentricity changes the detailed form of the perturbing potential from that in \\eref{Phi2}, but cannot change the precessional character leading to the torque equation (\\ref{align}). Our results remain unchanged, particularly the counteralignment condition (\\ref{cond}), apart from minor modifications of the timescale (\\ref{talign2}).  Thus in a random sequence of accretion events producing external discs, the prograde events have little effect, and the retrograde ones shrink the binary. In particular, a sequence of minor retrograde events with $J_{\\rm{d}} < J_{\\rm{b}}$ has a cumulative effect and must ultimately cause the binary to coalesce once the total retrograde $\\sum J_{\\rm{d}} = J_{\\rm{b}}$.  This is important, since the disc mass is limited by the onset of self--gravity to $M_{\\mathrm{d}} \\lesssim \\left(H/R \\right) M_{1}$ (cf King, Pringle and Hofmann 2008). Coalescence will then occur once the retrograde discs have brought in a total mass $M_2$, i.e. once a sequence of $\\gtrsim (M_2/M_1)(R/H)$ retrograde discs have accreted. For minor mergers this requires at most a few randomly oriented accretion disc events, rising to a few hundred for major mergers ($q >0.1$).  We note finally that similar considerations apply in planet--forming discs around stellar binary systems, which can also be initially misaligned \\citep{Bateetal2010}. This may offer a way of making retrograde planets in binaries, as recently suggested for $\\nu$ Octantis \\citep{EC2010}."
    },
    "1107/1107.2395_arXiv.txt": {
        "abstract": "By analysing models of the young massive cluster R136 in 30 Doradus, set-up using the herewith introduced and publicly made available code \\textsc{McLuster}, we investigate and compare different methods for detecting and quantifying mass segregation and substructure in non-seeing limited $N$-body data. For this purpose we generate star cluster models with different degrees of mass segregation and fractal substructure and analyse them.  We quantify mass segregation by measuring, from the projected 2d model data, the mass function slope in radial annuli, by looking for colour gradients in radial colour profiles, by measuring Allison's $\\Lambda$ parameter, and by determining the local stellar surface density around each star. We find that these methods for quantifying mass segregation often produce ambiguous results. Most reliable for detecting mass segregation is the mass function slope method, whereas the colour gradient method is the least practical in an R136-like configuration. The other two methods are more sensitive to low degrees of mass segregation but are computationally much more demanding. We also discuss the effect of binaries on these measures.  Moreover, we quantify substructure by looking at the projected radial stellar density profile, by comparing projected azimuthal stellar density profiles, and by determining Cartwright \\& Whitworth's $Q$ parameter. We find that only high degrees of substructure affect the projected radial density profile, whereas the projected azimuthal density profile is very sensitive to substructure. The $Q$ parameter is also sensitive to substructure but its absolute value shows a dependence on the radial density gradient of the cluster and is strongly influenced by binaries.  Thus, in terms of applicability and comparability for large sets of $N$-body data, the mass function slope method and the azimuthal density profile method seem to be the best choices for quantifying the degree of mass segregation and substructure, respectively. The other methods are computationally too demanding to be practically feasible for large data sets. ",
        "introduction": "\\label{Sec:Introduction} Understanding the process of star cluster formation is vital for astrophysics since most, if not all stars are born in a clustered mode \\citep{Lada03}. In the commonly accepted picture, star clusters form in three stages: first, a cold molecular cloud collapses and forms stars along filaments throughout this collapse. Secondly, the massive O- and B-stars start radiating off the residual gas until only a more or less bound ensemble of stars is left. Subsequently, the newly formed star cluster evolves dynamically until total dissolution (e.g. \\citealt{Kroupa01b, Portegies10}).  In this picture, the survival of star clusters throughout the birth process and the duration of the subsequent dynamical dissolution depends crucially on the star formation efficiency, i.e. how much of the cold gas gets transformed into stars. Moreover, it depends on the structure of the gas cloud and the distribution of the forming stars. That is, for a more massive ensemble of stars the gas expulsion process will be more violent, whereas for low-mass configurations gas expulsion will happen more adiabatically \\citep{Geyer01, Kroupa02b, Goodwin06, Bastian06, Baumgardt07, Baumgardt08b}. A further influence on the survival rate is given by the degree of mass segregation, i.e. when the most massive stars of an ensemble are located deeper within the forming cluster, they will have a more destructive influence on the subsequent evolution of the star cluster \\citep{Vesperini09}. A yet open question in this respect is whether or not star clusters form with primordial mass segregation, or if the initial mass function (IMF) of stars is the same throughout the whole star forming complex.  Observations of young clusters indeed suggest the existence of primordial mass segregation (e.g. \\citealt{Stolte02, Bontemps10b}). But such observed mass segregation is not necessarily due to variations of the IMF, as mass segregation can also develop quickly during the first few 100,000 years of cluster formation through dynamical relaxation. The timescale of this process is proportional to the relaxation time of the configuration and the mass ratio of the forming stars \\citep{Spitzer87}. Moreover, recent investigations show that initially substructured configurations can develop mass segregation on significantly shorter timescales \\citep{McMillan07, Allison09, Moeckel09, Allison10, Yu11}. In this picture, (fractal) substructures, which may have much shorter relaxation times, can segregate before they merge to form the final star cluster. Hence, substructure in young star clusters is not only a sign of a system not being virialised, but may also play a vital role in the process of star cluster formation. Thus, two major aspects of young star clusters have to be investigated in detail at the different stages of star cluster formation: the degree of mass segregation and the degree of substructure.  These questions may be addressed by means of collisional $N$-body computations. But such numerical investigations also require a choice of initial conditions. The question therefore remains what configuration star clusters have at the stage where dynamical investigations can set in. From hydrodynamical computations of collapsing gas clouds it appears that star formation may be happening in a hierarchical, fractal fashion \\citep{Klessen01, Bonnell03, Bonnell06}. During the subsequent dynamical evolution this substructure is erased and indeed a significant degree of mass segregation is established \\citep{Maschberger10}. But hydrodynamical computations only reach gas cloud masses of a few $10^3\\msun$, thus they do not shed any light on star formation in starburst regions or on the formation of globular clusters, nor can they account for the self-regulation induced by stellar feedback, yet.  Observations of young embedded star clusters show a similar picture like hydrodynamical computations. For example, \\citet{Lada03}, \\citet{Teixeira06}, \\citet{Allen07} as well as \\citet{Gennaro11} find many young clusters to show substructure and to be asymmetric. Moreover, \\citet{Gutermuth05, Gutermuth08} find young embedded clusters in near-IR data to be azimuthally asymmetric with a high degree of substructure. The star formation sites appear filamentary and elongated over scales of several parsec. Such observations suggest that young clusters expand, get more symmetric and lose substructure with ongoing gas removal \\citep{Gutermuth08, Bontemps10a}. The nearby, more evolved Orion Nebula Cluster, for example, shows a high degree of mass segregation which appears to be inconsistent with its current relaxation time \\citep{Hillenbrand98, Kroupa02a}. In this picture, this may indicate that the ONC also formed with a high degree of substructure.  But does this picture also apply to starbursts, and to the formation of globular cluster-like objects? Massive star formation sites with stellar masses above $10^4\\msun$ in the Milky Way like NGC 3603, Westerlund 1 and the Arches cluster, are rare and mostly heavily obscured by interstellar dust. Nevertheless, mass segregation as well as high degrees of substructure have been reported for those objects (e.g. \\citealt{Stolte02, Stolte06, Brandner08, Gennaro11}).  The most massive star formation site in the Local Group is the 30 Doradus complex in the Large Magellanic Cloud. It is the only nearby starburst region, which makes it the ``Rosetta Stone'' for understanding events such as the formation of globular clusters \\citep{Walborn91}. This paper therefore addresses this star formation site and aims at investigating mass segregation and substructure in this complex. For this purpose, we create models of the young massive cluster R136 which is forming in the 30 Doradus complex (Sec.~\\ref{Sec:R136} \\& \\ref{Sec:Models}). We then use various methods from the literature (Sec.~\\ref{Sec:Methods}) to detect and quantify mass segregation and substructure (Sec.~\\ref{Sec:Results}). Our aim is to test and calibrate these methods in order to be able to apply them to $N$-body computations which will be presented in a follow-up investigation, and to discuss their applicability to observational data. Furthermore, Appendix A contains a manual on our new star-cluster initialisation code \\textsc{McLuster}. ",
        "conclusions": "\\label{Sec:Conclusions} Here we introduce for the first time the code \\textsc{McLuster}, which is a publicly-available tool for initialising star cluster models and binary-rich stellar populations.  Moreover, we have tested and calibrated several methods for detecting and quantifying mass segregation and substructure in non-seeing limited star cluster data. We applied these methods to models of the young massive cluster R136 which is the only starburst site in the Local Group. Accessing this cluster and measuring its degree of mass segregation and substructure at its current age of about 3 Myr is of importance because it is the only object which can be resolved into single stars and at the same time can give insight to star formation conditions in starburst sites. Moreover, due to its size and mass we can assume that the present-day conditions in R136 can still yield insights to the conditions 3 Myr ago, as the dynamical time of this configuration is comparably long (about $0.3$ Myr), i.e. the cluster is dynamically young \\citep{Portegies10}. Finally, its moderate mass of about $10^5\\msun$ makes it accessible by means of $N$-body investigations. But for understanding the development of these two quantities, mass segregation and substructure, easily quantifiable and little time consuming methods have to be found.  We compare 4 different methods for quantifying mass segregation and 3 methods for quantifying substructure from the literature (Sec.~\\ref{Sec:Methods}). The former we quantify by comparing the mass function slope of massive stars, the radial colour profile \\citep{Gaburov08}, the minimum spanning tree measure \\citep{Allison09, Allison10} and the local surface density measure \\citep{Maschberger11}. We quantify substructure by looking at the projected radial number density profile, the projected azimuthal density profile \\citep{Gutermuth05, Gutermuth08} and by calculating the $Q$ parameter \\citep{Cartwright04}. For this purpose we set up star cluster models with different degrees of mass segregation and substructure using the new publicly available code \\textsc{McLuster} (Sec.~\\ref{Sec:Models}).  We find that the methods for detecting mass segregation often yield ambiguous results (Sec.~\\ref{Sec:Results}). From the four methods we compare, the mass function slope seems to be the simplest and most reliable measure for detecting mass segregation, even though it suffers from a somewhat arbitrary choice of radius, and significant detections of mass segregation are only possible for high degrees of mass segregation. The radial colour profile only yields significant results for the highest (and rather unrealistic) degrees of mass segregation for a configuration like R136. Both measures cannot handle substructure, i.e. clusters which are not radially symmetric. In contrast to that, the minimum spanning tree measure and the local surface density measure can handle substructured clusters \\citep{Maschberger11}, but their results are often ambiguous, and computationally they are much more demanding. Moreover, the minimum spanning tree method is strongly influenced by a high binary fraction. On the other hand, it has the unique advantage that the significance of its results is readily given. The local surface density measure is less influenced by binaries, and with careful calibration it can be a very sensitive method for detecting mass segregation. We recommend, though, that a measure of its significance similar to the standard deviation of the MST measure should be constructed.  For quantifying substructure we are left with the projected azimuthal density profile since the projected radial density profile only shows significant deviations from the spherical symmetric case for extremely substructured configurations. Such a cluster with a high degree of substructure can show strong bumps and kinks in its projected radial profile but those are difficult to quantify. The projected azimuthal density profile is a reliable measure of substructure (Fig.~\\ref{calibdADP}). We suggest to compute the mean projected azimuthal density and the normalised standard deviation from this mean to get a useful measure. The $Q$ parameter is also sensitive to substructure but is computationally very demanding, when using a standard algorithm like Prim's algorithm, as a minimum spanning tree for all cluster stars has to be constructed. If only a subset of stars is taken account for its computation, its absolute value shows a dependence on the number of stars in this subset which is due to the binary population we adopted. Thus, the $Q$ parameter is only of limited use in such configurations, unless a much more sophisticated algorithm like Kruskal's algorithm is used (see e.g. \\citealt{Lomax11}).  Finally, we have to conclude that most of the methods presented in this work will likely yield ambiguous results when applied to observations of young massive clusters, due to crowding and incompleteness effects, as has similarly been found by \\citet{Ascenso09}. Also the distance of such clusters and heavy, variable extinction will add further difficulties. But even for non-seeing limited data of (young) massive clusters, like $N$-body data, some of the methods will face great computational difficulties due to the large number of stars involved."
    },
    "1107/1107.1325_arXiv.txt": {
        "abstract": "We develop an analytic formalism that allows one to quantify the stability properties of \\xtype and \\ytype junctions in domain wall networks in two dimensions. A similar approach might be applicable to more general defect systems involving junctions that appear in a range of physical situations, for example, in the context of F- and D-type strings in string theory. We apply this formalism to a particular field theory, Carter's pentavac model, where the strength of the symmetry breaking is governed by the parameter $|\\epsilon|< 1$. We find that for low values of the symmetry breaking parameter \\xtype junctions will be stable, whereas for higher values an \\xtype junction  will separate into two \\ytype junctions. The critical angle separating the two regimes is given by  $\\qsubrm{\\alpha}{c} = 293^{\\circ}\\sqrt{|\\epsilon|}$ and this is confirmed using simple numerical experiments. We go on to simulate the pentavac model from random initial conditions and  we find that the fraction of \\xtype junctions to \\ytype junctions is higher for smaller $\\epsilon$, although \\xtype junctions do not appear to survive to late time. We also find that  for small $\\epsilon$  the evolution of the number of domain walls $\\qsubrm{N}{dw}$ in Minkowski space does not follow the standard $\\propto t^{-1}$ scaling law with the deviation from the standard lore being more pronounced as $\\epsilon$ is decreased. The presence of dissipation appears to restore the $t^{-1}$ lore. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.2921_arXiv.txt": {
        "abstract": "We demonstrate estimating the total infrared luminosity, \\LIR, and star formation rates (SFRs) of star-forming galaxies at redshift $0 < z < 2.8$ from single-band 24 $\\micron$ observations, using local SED templates without introducing additional free parameters. Our method is based on characterizing the spectral energy distributions (SEDs) of galaxies as a function of their \\LIR\\ surface density, which is motivated by the indications that the majority of IR luminous star-forming galaxies at $1 < z < 3$ have extended star-forming regions, in contrast to the strongly nuclear concentrated, merger-induced starbursts in local luminous and ultraluminous IR galaxies. We validate our procedure for estimating \\LIR\\ by comparing the resulting \\LIR\\ with those measured from far-IR observations, such as those from {\\it Herschel} in the Extended {\\it Chandra} Deep Field South (ECDFS) and {\\it Hubble} Deep Field North (HDFN), as well as \\LIR\\ measured from stacked far-IR observations at redshift $0 < z < 2.8$. AGNs were excluded using X-ray and $3.6-8.0$ $\\micron$ observations, which are generally available in deep cosmological survey fields. The Gaussian fits to the distribution of the discrepancies between the \\LIR\\ measurements from single-band 24 $\\micron$ and {\\it Herschel} observations in the ECDFS and HDFN samples have $\\sigma < 0.1$ dex, with $\\sim 10\\%$ of objects disagreeing by more than 0.2 dex. Since the 24 $\\micron$ estimates are based on SEDs for extended galaxies, this agreement suggests that $\\sim 90\\%$ of IR  galaxies at high $z$ are indeed much more physically extended than local counterparts of similar \\LIR, consistent with recent independent studies of the fractions of galaxies forming stars in the {\\it main-sequence} and {\\it starburst} modes, respectively. Because we have not introduced empirical corrections to enhance these estimates, in principle, our method should be applicable to lower luminosity galaxies. This will enable use of the 21 $\\micron$ band of the Mid-Infrared Instrument (MIRI) on board the {\\it JWST} to provide an extremely sensitive tracer of obscured SFR in individual star-forming galaxies across the peak of the cosmic star formation history. ",
        "introduction": "The mid-infrared (mid-IR) is a unique window to study the evolution of star-forming galaxies, especially at redshift $1 < z < 3$ where the cosmic star formation rate (SFR) peaks \\citep[e.g.,][]{HopkinsBeacom06}. At these redshifts, a majority of star formation took place in obscured environments, where dust reprocesses the UV photons from hot young stars into IR emission \\citep[see e.g.,][]{LeFloch05, PPG05, Dole06, Buat07}. Star-forming galaxies at these redshifts also exhibit a large spread of extinction values and diverse dust distribution scenarios \\citep{Rujopakarn12}. These factors pose inherent challenges for optical and UV estimators of the SFR, which need to be complemented by IR techniques.  In the past decade, {\\it ISO}, {\\it Spitzer}, {\\it WISE}, {\\it AKARI}, and {\\it Herschel} have allowed us to study star formation from the local Universe out to high $z$ using mid-IR and far-IR observations. Measurements that determine the IR luminosity, \\LIR, from these missions trace the energy absorbed from UV photons emitted by short-lived massive stars \\citep[see, e.g.,][]{Kennicutt98}. The most direct approach to measure \\LIR\\ is to use multi-band mid- and far-IR photometry. Recent examples use {\\it Herschel} to observe star-forming galaxies at $100 - 500$ $\\micron$ and fit galaxy SEDs to measure \\LIR\\ \\citep[e.g.,][]{Elbaz10, Rex10, Elbaz11}. The complete characterization of the spectral peak of the dust emission provides a good measurement of \\LIR\\ and SFR. However, the inherent requirement of multi-band photometry for far-IR SED fitting compromises this approach for faint galaxies whose detection is limited by confusion noise, particularly at the longer wavelengths \\citep[e.g.,][]{Condon74, Dole04}. A second approach is to use a monochromatic (single-band) IR luminosity to trace \\LIR\\ and SFR. Locally, the rest-frame single-band 24 $\\micron$ luminosity has been shown to be one of the best tracers for \\LIR\\ and SFR \\citep{Calzetti07, Rieke09}. Although the longer wavelengths, such as 70 $\\micron$ and 160 $\\micron$, are closer to the peak of the thermal dust emission and thus reduce the size of the bolometric corrections, they are affected more by the cold dust heated by old stars, which increases scatter in the SFR calibration compared to estimates from 24 $\\micron$ \\citep{Rieke09, Kennicutt09}, particularly in less luminous IR galaxies \\citep{Calzetti10a}.  Beyond the local Universe, however, the redshifted 24 $\\micron$ band probes wavelengths containing the aromatic emissions (e.g., polycyclic aromatic hydrocarbons, hereafter PAH); the strongest PAH emission complexes at 6.2, 7.7, and 8.6 $\\micron$ \\citep[e.g.,][]{JDSmith07} redshift into the 24 $\\micron$ band by $z \\sim 2$. Although they help boost the 24 $\\micron$ flux and aid detection of galaxies at high $z$, PAHs pose two challenges to using single-band 24 $\\micron$ observations to estimate \\LIR\\ and SFR. First, the PAH emission is influenced by environment (e.g., UV radiation, optical depth), introducing a $\\sim 0.2-$dex scatter to the \\LIR\\ and SFR estimates \\citep{Roussel01}. Second, the PAH emission appears to strengthen intrinsically at high $z$ compared to that found in local galaxies with the same \\LIR\\ \\citep{Rigby08, Farrah08, Takagi10}. Therefore, the bolometric corrections to determine \\LIR\\ measured from local galaxies in the PAH wavelength  regions will not be applicable at high $z$. Recent studies  indicate that applying the local bolometric corrections to high-$z$ galaxies will overestimate their \\LIR\\ and SFR by up to an order of magnitude. This is reported as the ``mid-IR excess'' in recent far-IR studies using {\\it Spitzer} and {\\it Herschel} \\citep[e.g.,][]{Papovich07, Elbaz10, Nordon10, Barro11, Nordon11}. The SED evolution causes a {\\it systematic bias} that must be taken into account to use 24 $\\micron$ observations as \\LIR\\ and SFR indicators beyond the local Universe.  Out to $z \\sim 3$, the {\\it Spitzer} 24 $\\micron$ observations probe weaker SFR at any given redshift than is possible with the far-IR bands \\citep[see Figure 4 of][]{Elbaz11}. For {\\it JWST}, 21 $\\micron$ is the longest wavelength band suitable for deep cosmological surveys. Therefore, our understanding of star-forming  galaxies at high $z$ will depend critically on mid-IR SFR indicators, where the current state-of-the-art prescription to estimate \\LIR\\ from single-band 24 $\\micron$ observations still presents a $0.4-$dex scatter \\citep{Nordon11}. Additionally, the {\\it Spitzer} 24 $\\micron$ data are already available from deep legacy surveys (e.g., GOODS, FIDEL, COSMOS, SpUDS), as well as from large-area surveys (e.g., SWIRE and the Bo\\\"{o}tes field) that will not be fully surveyed by current facilities to the same depth in terms of SFR. Exploration of means to reduce the current $0.4-$dex scatter in \\LIR\\ estimates will allow utilization of the mid-IR observations to the fullest extent in the upcoming decade.  In this paper, we apply the results from our previous study, \\citet{Rujopakarn11}, to take into account the SED evolution of star-forming galaxies and refine the 24 $\\micron$ \\LIR\\ and SFR indicators. \\citet{Rujopakarn11} find that the IR luminosity surface density, \\LIRSD, affords an accurate description of the SED and subsequently allows accurate bolometric corrections out to high $z$, specifically out to $z = 2.8$, the farthest redshift where the {\\it Spitzer} 24 $\\micron$ band traces predominantly PAH emissions. The measurement of \\LIRSD\\ requires high-resolution imaging of the star-forming regions, which is only available for a small number of galaxies. We use these galaxies as a tool to construct a simple formula to estimate \\LIR\\ and SFR without measuring \\LIRSD\\ for individual galaxies and using only single-band 24 $\\micron$ flux and redshift measurements, and then compare the resulting \\LIR\\ estimates with results using far-IR data.  This paper is organized as follows. We discuss the evolution of the SEDs of star-forming galaxies and the use of \\LIRSD\\ as a guide to select the appropriate SED for high-$z$ galaxies in Section \\ref{sec:sedevol}. The derivation of the single-band \\LIR\\ indicator is described in Section \\ref{sec:indicator}. We then compare the new \\LIR\\ indicator with other independent measurements in Section \\ref{sec:indicator_test} and present the \\LIR\\ SFR indicator in Section \\ref{sec:final_SFR}. Lastly, we discuss the validity of rest-frame $8-24$ $\\micron$ as a tracer of \\LIR\\ and SFR, along with implications of our results in Section \\ref{sec:discuss}. Throughout this paper, we assume a $\\Lambda$CDM cosmology with  $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70$~km~s$^{-1}$Mpc$^{-1}$. We follow the convention of \\citet{Rieke09} and adopt the definition of \\LIR\\ of \\citet{SandersMirabel96}; these studies defined \\LIR\\ slightly differently (i.e., $5-1000$ $\\micron$ vs. $8 - 1000$ $\\micron$), but the resulting \\LIR\\ values are closely consistent.  We will refer to galaxies with \\LIR\\ in the range of $10^{11}-10^{12}$ and those with \\LIR\\ $>10^{12}$ \\Lsun\\ as Luminous Infrared Galaxies (LIRGs) and Ultraluminous Infrared Galaxies (ULIRGs), respectively, or collectively as U/LIRGs.\\\\ ",
        "conclusions": "\\label{sec:conclusions} The use of single-band mid-IR indicators (such as 24 $\\micron$ observations) to estimate \\LIR\\ has previously been affected by the differences between the structure of local U/LIRGs, upon which the SED templates were constructed, and those of the galaxies in the cosmological surveys being studied. Starbursts in the local U/LIRGs are interaction-induced and compact (sub-kpc), while U/LIRGs at high $z$ are typically much more extended ($\\sim$ few kpc), with surface areas $\\sim 100\\times$ larger than their local counterparts. The resulting larger surface area of the photodissociation regions harboring aromatic emission leads to stronger aromatic features at high $z$ that appears as an evolution in aromatic strength with $z$. The ultimate implication of this morphological difference is that single-band mid-IR observations, which at $z > 1$ probe the aromatic-dominated SED region, will overpredict \\LIR\\ and SFR compared to far-IR measurements if the relationship of local SED templates to \\LIR\\ is assumed for the calculation of the bolometric and k$-$corrections.  Following the \\citet{Rujopakarn11} result that \\LIRSD\\ can serve as a good predictor for the appropriate SED of star-forming galaxies, we construct a new 24 $\\micron$ single-band \\LIR\\ and SFR indicator for star-forming galaxies at redshift $0 < z < 2.8$. The resulting $0.03$ dex average agreement (0.13 dex scatter) between \\LIR\\ estimates from the new single-band 24 $\\micron$ indicator and those from multi-band far-IR SED fitting support the \\citet{Rujopakarn11} result that \\LIRSD\\ is indeed the dominant factor affecting the SED shapes of star-forming galaxies.  For the purpose of estimating \\LIR\\ and SFR, we recommend using the new indicator at $0 < z < 2.8$. A step-by-step recipe is given in Section \\ref{sec:final_SFR}. The new indicator converges with that of \\citet{Rieke09} for local and low $z$ galaxies (Section \\ref{sec:discuss_sfmodes}). But unlike the \\citet{Rieke09} indicator, the new indicator carries the extended structure assumption and could underestimate \\LIR\\ for the minority of cases ($10\\%$) of high-$z$ starbursts that are compact (Figures \\ref{compare_lir_CDFS} and \\ref{compare_lir_HDFN}). Thus, if an object is known {\\it a priori} to be a compact starburst (e.g., from high-resolution optical or interferometric radio observations), we recommend readers to apply the bolometric correction factor given by \\citet{Rieke09} in their Equation 14. On the other hand, where the structure of the star-forming region is not known {\\it a priori}, we quote an overall scatter of 0.13 dex, when AGNs are identified and excluded from the sample using X-ray and $3.6-8.0$ $\\micron$ (i.e. {\\it Spitzer}/IRAC) observations.  We thank Daniel Eisenstein for insightful discussions, and J. D. Smith for the L$_{{\\rm PAH}}/$\\LIR\\ data. This work is supported by contract 1255094 from Caltech/JPL to the University of Arizona. WR gratefully acknowledges the support from the Thai Government Scholarship.  \\appendix"
    },
    "1107/1107.0862_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) and supernovae (SNe) bring new perspectives to the study of neutron stars and white dwarfs, as well as opening new branches of theoretical physics and astrophysics. ",
        "introduction": "\\label{sec:intro}  I dedicate this talk to John Archibald Wheeler recalling some crucial moments in our collaboration and focusing on two of the most energetic and transient phenomena in the universe: supernovae (SNe) and gamma-ray bursts (GRBs). From the knowledge being gained daily on these astrophysical systems a new physical and astrophysical understanding is emerging. The physics of neutron stars and black holes is rapidly evolving, new more complex astrophysical scenarios previously unimaginable are being developed and new domains of fundamental physics and astrophysics are being opened to scrutiny.  I will also outline some of the ideas I advanced with Wheeler in our book ``Black Holes Gravitational Waves and Cosmology'' \\cite{1974bhgw.book.....R} and in our article ``Introducing the Black Hole'' \\cite{1971PhT....24a..30R} (see Figs.~\\ref{fig_ibh}, \\ref{fig_irrmass}, \\ref{fig_uniqueness}, \\ref{effpot}): they are coming to full fruition and new additional topics need urgent attention.   \\begin{figure}[t] \\centering \\includegraphics[width=0.70\\hsize,clip]{introbh} \\includegraphics[width=0.70\\hsize,clip]{ergo} \\caption{\\textbf{Above:} It was in the writing of our ESRO report \\cite{1971ESRSP..52...45R} with Johnny, which later became the first chapters of our book on ``Black Holes, Gravitational Waves and Cosmology'' \\cite{1974bhgw.book.....R} that some of the crucial new ideas on the last phases of gravitational collapse leading to a black hole were formulated. They were summarized in the celebrated article ``Introducing the Black Hole'' \\cite{1971PhT....24a..30R}. There, the crucial new concept of a ``black hole'', as a physical and astrophysical system and not just as an analytic solution of the Einstein equations, was presented for the first time. \\textbf{Below:} The particle decay process in the field of a black hole: figure and original caption reproduced from Ref.~\\refcite{1970PhRvL..25.1596C}.} \\label{fig_ibh} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[width=0.59\\hsize,clip]{ergo_bw} \\caption{The concept of the ``ergosphere'' was introduced \\cite{1971PhT....24a..30R} as well as the concept of extraction of rotational and electromagnetic energy from a black hole with their corresponding thermodynamical analogies, further formalized in Refs.~\\refcite{1970PhRvL..25.1596C,1971PhRvD...4.3552C}.} \\label{fig_irrmass} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[width=\\hsize,clip]{BH_Uniqueness_1} \\caption{Wheeler described another crucial property of black holes in a very provocative pictorial representation illustrating the concept of the black hole uniqueness, namely the unprecedented situation in physics that the final output of the process of gravitational collapse was only dependent on three parameters, mass, charge, and angular momentum, independent of the original details of the precursor astrophysical source. We shall see shortly how this uniqueness concept plays a fundamental role in GRBs.} \\label{fig_uniqueness} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[width=0.7\\hsize,clip]{effpot} \\includegraphics[width=0.7\\hsize,clip]{bind} \\caption{\\textbf{Above:} The effective potential experienced by a test particle moving in the equatorial plane of an extreme Kerr black hole. For corotating orbits, with positive values of the angular momentum, the maximum binding of $42.35$\\% of the rest mass of the test particle is reached at the horizon. Details in Refs.~\\refcite{1974bhgw.book.....R,1975ctf..book.....L}. \\textbf{Below:} Nuclear vs.\\ gravitational binding energy in a Schwarzschild black hole compared and contrasted. The gravitational binding energy in the Kerr case is even bigger. See Ref.~\\refcite{1974bhgw.book.....R}, also quoted in Ref.~\\refcite{1975ctf..book.....L}.} \\label{effpot} \\end{figure}  It has been traditionally established that the understanding of an astrophysical phenomenon is only reached when its energetic aspects have been properly identified. This has certainly occurred in three cases of clear success in the understanding of: \\begin{enumerate} \\item the stellar evolution, which was reached after thermonuclear reactions were fully tested in Earth-bound experiments and their results applied to astrophysics (see Refs.~\\refcite{p20,e20,1929ZPhy...52..496G,1929ZPhy...54..656A,vw37,vw38,1939PhRv...55..103B,1957RvMP...29..547B}); \\item pulsars, reached as soon as the rotational energy of neutron stars was recognized as the energy source of these phenomena (see Refs.~\\refcite{1966ARA&A...4..393W,1969ApJ...155L.107F,1968Natur.217..709H}); \\item the nature of binary X-ray sources, achieved as soon as the role of the gravitational energy in the fully general relativistic regime was established; the already well known results for a Schwarzschild metric were extended to the Kerr metric (see Fig.~\\ref{effpot}, Ref.~\\refcite{1974bhgw.book.....R}, also quoted in Ref.~\\refcite{1975ctf..book.....L}, and Refs.~\\refcite{1978pans.proc.....G,2003IJMPA..18.3127G}); these results were applied to the study of neutron stars and black holes in binary X-ray sources (see Ref.~\\refcite{1974asgr.proc..349R}). \\end{enumerate}  If we turn from these very well understood and long duration phenomena to the most energetic and transient ones such as SNe and GRBs, it is well accepted that gravitational energy in the general relativistic regime plays the principal role in their explanation, namely the energy originating in the process of gravitational collapse either to a neutron star or to a black hole. Such a  release of gravitational energy in supernovae is necessary to trigger the release of thermonuclear energy processes as well as the copious emission of neutrinos leading to the formation of a neutron star. Similarly, the more extreme general relativistic conditions created in the collapse to a black hole trigger the vacuum polarization quantum process with a vast production of electron-positron pairs leading to the creation of a GRB. Nevertheless, some open issues remain in both cases.  \\subsection{Some open issues on SNe}  For the case of SNe there is no better prototype then the Crab Nebula see Fig.~\\ref{fig_crab}.  \\begin{figure}[t] \\centering \\includegraphics[width=\\hsize,clip]{crab} \\caption{The Crab Nebula has certainly been one of the most complex systems ever analyzed by the human mind. Introduced by Messier in his famous catalog of nebulae \\cite{messier} as M1, it was discovered by John Bevis and the name ``Crab Nebula'' was attributed by William Parsons, 3rd Earl of Rosse. The first images of its filamentary structure were obtained by Baade \\cite{1942ApJ....96..188B}. It was thanks to the fortunate interaction between Jaan Oort \\cite{1942PASP...54...95M} and the sinologist J.J. Duyvendak \\cite{1942PASP...54...91D} that it was possible to identify the Crab nebula as the remnant of the SN which exploded in 1054 A.D. This event was recorded in the Chinese chronicles of the Sung dynasty and in the Japanese records Mei-Getsuki. This fascinating history has been reconstructed in the book by I.S. Shklovsky on SNe \\cite{1968QB895.S6213....}.} \\label{fig_crab} \\end{figure}  The preliminary understanding of the role of a neutron star in the process of gravitational collapse leading to a supernova and the fundamental role of supernovae in the creation of cosmic rays were presented  in a short and epochal paper by Baade and Zwicky \\cite{1934PNAS...20..254B}. Following the work of George Gamow \\cite{g38}, Oppenheimer and Volkoff \\cite{1939PhRv...55..374O} gave the first general relativistic computations for neutron stars and introduced the existence of a  critical mass against gravitational collapse. The continued gravitational collapse, solely described by general relativity for a star  with mass larger than the critical mass, was then developed by Oppenheimer and Snyder \\cite{1939PhRv...56..455O}. Gamow and Schoenberg \\cite{1941PhRv...59..539G} identified the  fundamental role of the neutrino and antineutrino emission in order to dissipate the enormous thermal energy developed in the early phases of gravitational collapse. The neutrino-antineutrino emission, through the URCA process, introduced the essential cooling needed for the formation of a neutron star.  In a pioneering work, Hoyle and Fowler \\cite{1960ApJ...132..565H} described the more complex thermonuclear process occurring in the gravitational collapse to a neutron star. They clearly differentiate two different possibilities: the gravitational collapse starting from a white dwarf, with a mass larger than the Stoner-Chandrasekhar critical mass \\cite{2010arXiv1012.0154R}, leading to a type I SN; or, alternatively, the gravitational collapse starting from the larger core of a more massive star, leading to a type II SN. The very extensive work on thermonuclear processes preceding and occurring at the onset of gravitational collapse has been exhaustively summarized in the classic book of David Arnett \\cite{1996sunu.book.....A}. The observations of the supernova 1987 A together with the associated neutrinos was a particularly important step in the verification of this very complex scenario \\cite{1989ARA&A..27..629A}.  There are nevertheless still open questions about the release of the gravitational energy: still unexplained is the process of the expulsion of the SN remnant. We still lack an understanding of such a fundamental process. The possibility of explaining such an expulsion process by the very large flux of neutrinos emitted during the cooling of the collapsing core has been advocated by Jim Wilson and his collaborators \\cite{1985ApJ...295...14B,1985nuas.conf..422W,1986NYASA.470..267W}. But it is a matter of fact that still today the theory is  unable to explain this expulsion process \\cite{1996sunu.book.....A}.  My collaborators and I, working in the field of relativistic astrophysics, are currently focused on seeing whether the way out of this impasse might be found in some crucial properties of neutron stars which have been neglected so far in a oversimplified approach. This crucial idea may well also have been missed in the extremely complex and time consuming numerical simulations.  There is no doubt that there is plenty of gravitational energy available in the process of gravitational collapse. The gravitational energy can be estimated simply by (see e.g.\\ Ref.~\\refcite{1967aits.book.....C} and references therein): \\begin{equation} E = \\frac{3}{5}\\frac{GM^2}{R}\\,, \\label{eg1} \\end{equation} where $M$ and $R$ are the mass and radius of the neutron star. If one assumes for the final radius of the collapsing star \\begin{equation} R \\sim \\alpha \\frac{GM}{c^2}\\,, \\label{eq_2} \\end{equation} then the gravitational energy is \\begin{equation} E \\sim \\frac{1}{\\alpha} Mc^2\\,. \\label{eg2} \\end{equation} For $10 \\lesssim \\alpha \\lesssim 10^2$, typical for a neutron star, one readily obtains gravitational energies on the order of $10^{52}$--$10^{53}$ erg, which can well explain the observed energy of SNe in the range $10^{49}$--$10^{51}$ erg. In order to utilize this gravitational energy, either for the expulsion of the remnant or for the observed electromagnetic or particle energy flux emitted during the process of gravitational collapse, it appears essential to \\emph{minimize} the kinetic energy of the implosion of the core and transform the gravitational energy into an alternative form of energy with the necessary explosive power. There are two obvious mechanisms for obtaining this minimization process for the kinetic energy of the implosion: either the presence of rotation or the presence of new electrodynamical processes. The role of rotational energy is clearly revealed by the rotational energy of the newly formed neutron star, which is then released on time scales of thousands to millions of years in pulsars.  Turning to alternative instantaneous braking mechanisms occurring during the process of gravitational collapse, we are currently exploring the possibility of using electrodynamical processes. In our analysis, even the electrodynamics of the neutron star configurations of equilibrium, especially in the interface of the crust to the core of the neutron star, needs further scrutiny. Until now the condition of local charge neutrality in the neutron star configuration of equilibrium has been adopted mainly for mathematical simplicity (see Sec.~\\ref{sec:loc_glob}). The current stringent energy release requirements, both in SN events and in GRBs, as we will see shortly, demand a deeper analysis of the electrodynamical properties during the gravitational collapse process. A good starting point is the study of the equilibrium of a neutron star imposing global but not local charge neutrality. In such cores overcritical electric fields can develop and a large amount of electron-positron plasma can in principle be created during the process of their gravitational collapse, leading to explosive effects. This field of research will be summarized briefly in the conclusions.  \\begin{figure}[t] \\centering \\includegraphics[width=0.49\\hsize,clip]{vela1} \\includegraphics[width=0.49\\hsize,clip]{vela_firstburst} \\caption{\\textbf{Left:} the Vela 5A and 5B satellites. \\textbf{Right:} a typical event as recorded by three of the Vela satellites. Details in Ref.~\\refcite{1975ASSL...48.....G} by I. Strong.} \\label{fig_vela} \\end{figure}  \\subsection{GRBs: The first steps of the fireshell vs.\\ fireball scenario and the three paradigms}  For introducing the topic of GRBs there is no better illustration of this momentous research then the images of the original Vela Satellites (see Fig.~\\ref{fig_vela}). The discovery of GRBs was a fortunate outcome of one of the most alarming episodes of the cold war between the United States of America and the then Soviet Union. The mission was inspired by an unconventional proposal by our colleague Yacob Borisovich Zel'dovich \\cite{2010AIPC.1205....1R,ruKerr,2009AIPC.1132..199R}. The first public announcement of the GRB discovery occurred in a session I organized with Herb Gursky at the AAAS meeting in San Francisco \\cite{1975ASSL...48...47S}. On the meager observational evidence of the results presented in San Francisco, I worked out a model for GRBs with Thibault Damour, based on the possibility of explaining the origin of their gamma and X-ray fluxes by the process of vacuum polarization ``a la Heisenberg-Euler-Schwinger'' occurring around a black hole endowed with electromagnetic structure \\cite{1975PhRvL..35..463D}. In that paper, written a few months after the announcement of the GRB discovery, we clearly pointed out three aspects: \\begin{enumerate} \\item The role of the extractable electromagnetic energy of the black hole, or, more simply, the ``blackholic'' energy \\cite{ruKerr}, as the GRB energy source, implicit in the mass-energy formula of the black hole given by Christodoulou and Ruffini \\cite{1971PhRvD...4.3552C}: \\begin{equation} \\left\\{ \\begin{array}{l} m^2 = \\left(m_{ir} + \\frac{e^2}{4m_{ir}}\\right)^2 + \\frac{L^2}{4m_{ir}^2}\\, ,\\\\[6pt] S = 16 \\pi m_{ir}^2\\, ,\\\\[6pt] \\frac{L^2}{4m_{ir}^4} + \\frac{e^4}{16m_{ir}^4} \\leq 1\\, ,\\\\[6pt] \\delta S = 32 \\pi m_{ir} \\delta m_{ir} \\ge 0\\, , \\end{array} \\right. \\label{irrmass} \\end{equation} where $m$ is the total mass-energy of the black hole, $m_{ir}$ the irreducible mass, $S$ the surface area and $e$ and $L$ are the black hole charge and angular momentum in geometrical units ($G=c=1$). From this formula it follows that the amount of energy stored in a black hole, and in principle extractable by reversible transformations during the process of gravitational collapse, could be as high as $29\\%$ of the rotational energy and $50\\%$ of the electromagnetic energy of the black hole. The existence of what became known later as the ``blackholic energy'' extraction process (see Ref.~\\refcite{ruKerr}) had its beginnings here. In order to have an instantaneous explosive process, we directed our attention to the electromagnetic energy component in the mass-energy formula. We expected that the extraction of rotational energy, the other form of the blackholic energy, would take place on a much longer time scale, on the order of millions of years, e.g.\\ in active galactic nuclei. \\item The creation of an electron-positron pair plasma by the vacuum polarization process in the field of a Kerr-Newman black hole, representing the actual process of energy extraction from the black hole; these processes of $e^+e^-$ pair creation do approach the reversible transformations of a black hole introduced in Ref.~\\refcite{1971PhRvD...4.3552C}. \\item The identification of a typical ``energy scale'' on the order of $\\sim 10^{54} M_{BH}/M_\\odot$ erg, where $M_{BH}$ is the black hole mass, as the typical energy of a GRB to be expected in our model. \\end{enumerate} The theoretical background of this work was well grounded: \\begin{itemize} \\item on the mathematical side the solution of the Einstein-Maxwell equations found by Roy Kerr \\cite{1963PhRvL..11..237K}, and by Ted Newman and collaborators \\cite{1965JMP.....6..918N}, as well as, on the physical side, the mass-energy formula of the Kerr-Newman black holes (see Eq.~\\ref{irrmass}); \\item on the quantum physics side, the Heisenberg-Euler \\cite{1936ZPhy...98..714H} and Schwinger \\cite{1951PhRv...82..664S} formalism for the creation of an $e^+e^-$ pair plasma by the vacuum polarization process; the basic equations in general relativity have been properly identified in Ref.~\\refcite{1975PhRvL..35..463D} and fully reviewed in Ref.~\\refcite{2010PhR...487....1R}; \\item on the observational side, after the discovery by the Vela satellites a large momentum was gained in the international community making GRBs a main object of the observations of a large variety of dedicated missions worldwide. One of the major issues addressed at the time was to know the actual location and distance of these sources in the universe. Just knowing their flux on the Earth, in the absence of such a distance determination, it was not possible to know their energetics. The GRB energies, in fact, could vary over an enormous range of values depending on their location in the solar system, in the galaxy, or in the distant extragalactic universe. Consequently, an enormous number of models existed \\cite{ruKl}. This problem was finally overcome by the epochal observations by the BeppoSAX satellite, unequivocally confirming the typical ``energy scale'' I had previously introduced with T. Damour\\cite{1997Natur.387..783C}. \\end{itemize}  The Compton Gamma-Ray Observatory satellite, with the eight BATSE detectors, was the first of many very successful missions devoted to the study of GRBs. The outcome of these early observations led to two remarkable discoveries: the substantial isotropy of the distribution of the sources in the sky (see Fig.~\\ref{fig_isotr}), as well as a clear separation of two families of events: the short GRBs, with observed duration shorter than $\\sim 1$ seconds, and the long GRBs, corresponding to the complementary case (see Fig.~\\ref{fig_tavani}).  \\begin{figure}[t] \\centering \\includegraphics[width=0.49\\hsize,clip]{cgro_foto} \\includegraphics[height=0.49\\hsize,angle=90,clip]{batse2k} \\caption{The Compton Gamma-Ray Observatory satellite and the position in the sky of the observed GRBs in galactic coordinate. Different colors correspond to different intensities at the detector. There is almost perfect isotropy, both in the spatial and in the energetic distributions.} \\label{fig_isotr} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[width=\\hsize,clip]{tavani} \\caption{The energy fluence-averaged hardness ratio for short ($T < 1$) and long ($T> 1$ s) GRBs are represented. Reproduced with the kind permission of M. Tavani, from Ref.~\\refcite{1998ApJ...497L..21T} where the details are given.} \\label{fig_tavani} \\end{figure}  \\begin{figure}[t] \\centering \\includegraphics[width=\\hsize,clip]{satelliti} \\caption{The large flotilla of space observatories which followed the Vela satellites is here represented by the sequence of the Compton satellite and the fundamental contribution by the BeppoSAX satellite which, through its discovery of the afterglow, has allowed the determination of a sharper position for GRBs and consequently the possibility to identify their optical and radio counterparts with many space and Earth-based telescopes. This has allowed the determination of their cosmological distance and of their isotropic equivalent energy on the order of $10^{54}$ erg. Currently, the satellites XMM, Chandra, Swift, Fermi and Agile are giving new fundamental contributions on the energy, time variability and spectra of GRBs.} \\label{fig_satelliti} \\end{figure}  The number of dedicated missions soon increased (see Fig.~\\ref{fig_satelliti}). The crucial discovery came from the Italian-Dutch satellite BeppoSAX \\cite{1997Natur.387..783C}, jointly operating, for the first time, in both the X-ray and gamma-ray domain. This allowed the ``prompt radiation'' observed in gamma-rays to be associated with a long-lasting ($10^6$--$10^7$ s) ``afterglow'' observed in X-rays. This extremely important discovery has led to a better understanding of the nature of the GRB sources and offered the essential tool for improving their localization in the sky. It triggered a coordinated and complementary set of observations in the optical and radio energy bands: in turn, the optical identification, made possible by the timely development of the most powerful family of optical telescopes at Mauna Kea and at the VLT in Chile, allowed the determination of the GRB distance and therefore of their energetics (see Fig.~\\ref{fig_satelliti}). This led to distances of the sources with $z$ in the range between $z \\sim 0.0084$ and $z \\sim 8.2$ \\cite{2009ARA&A..47..567G,2009Natur.461.1254T,2009Natur.461.1258S}, and corresponding values of the energetics all the way up to the order of $10^{54}$ erg, exactly in the range we had predicted with Thibault Damour 25 years earlier.  We soon returned to our model by introducing three major new conceptual developments: \\begin{enumerate} \\item The identification of the region around the black hole where an overcritical electric field could occur, leading to the creation of $e^+e^-$ pairs out of the vacuum. The $e^+e^-$ were assumed to be in thermal equilibrium; the first identification was performed around a Reissner-Nordstr{\\o}m black hole. From the Greek name ``$\\delta\\acute{\\upsilon}\\alpha\\varsigma$, $\\delta\\acute{\\upsilon}\\alpha\\delta{o}\\varsigma$'' for ``pairs'', the name ``dyadosphere'' was introduced \\cite{1998bhhe.conf..167R,1998A&A...338L..87P} for this region, characterized by a total energy $E_{tot}^{e^\\pm}$, a mean energy of the pair created on the order of $1$--$4$ MeV and a characteristic radius $10^8$--$10^9$ cm. For details see Ref.~\\refcite{1998A&A...338L..87P}. \\item The study with Jim Wilson of the dynamics of the $e^+e^-$ pairs as the fundamental acceleration process of GRBs, reaching bulk Lorentz gamma factors $\\gamma \\sim 10^2$--$10^3$ (the pair-electromagnetic pulse, PEM pulse). This theoretical result was obtained in Ref.~\\refcite{1999A&A...350..334R} (for details see Sec.~\\ref{sec_canonical}). \\item Still with Jim Wilson, the study of the acceleration process of the optically thick $e^+e^-$-baryon plasma (the pair-electromagnetic-baryon pulse, PEMB pulse); this theoretical result was published in Ref.~\\refcite{2000A&A...359..855R}. This analysis introduced the new concept of plasma baryon loading defined by $B=M_{B}c^{2}/E_{tot}^{e^\\pm}$, where $M_B$ is the total mass of the baryons. To our surprise, we found a maximum value of the baryon loading for the existence of the electron-positron acceleration process $B \\leq 10^{-2}$ (for details see Sec.~\\ref{sec_canonical}). \\end{enumerate}  Some of these considerations have been consistent with and others in contrast to other articles in the contemporary literature: \\begin{enumerate} \\item The role of the $e^+e^-$ plasma, originally presented in our work with Damour \\cite{1975PhRvL..35..463D}, was soon generally adopted, with and without references! \\cite{1986ApJ...308L..47G,1990ApJ...365L..55S,1993ApJ...405..278M,1993ApJ...415..181M,1993MNRAS.263..861P,1994ApJ...430L..93R}; \\item There has been a substantial difference between our approach, purporting an initial thermal distribution of the $e^+e^-$ plasma in the dyadosphere, and the work by Cavallo and Rees \\cite{1978MNRAS.183..359C}, purporting the total annihilation of the $e^+e^-$ pairs originating in a process of gravitational collapse; \\item The model ``a la'' Cavallo-Rees has led to the concept of a ``fireball,'' a very hot cavity originating from the annihilation of the $e^+e^-$ pairs pushing on the surrounding circumstellar matter. In contrast, our model has led to the alternative concept of the ``fireshell'' in which the gradual annihilation of the $e^+e^-$ pairs lead to a self-acceleration of an optically thick shell \\cite{1999A&A...350..334R,2000A&A...359..855R}. The cavity inside such a shell is practically at zero temperature, and the $e^+e^-$ pairs gradually annihilate lasting until the fireshell becomes optically thin \\cite{2000A&A...359..855R}. \\end{enumerate}  To make this general situation even more exciting, a totally unexpected phenomena has been observed: the discovery of a GRB family that is relatively weak ($10^{49}$--$10^{51}$ erg) and close ($z < 0.17$), coincident in space and time with a SN event (see e.g.\\ Ref.~\\refcite{Mosca_Orale} and references therein).  The ``majority'' point of view reached consensus on three main points: \\renewcommand{\\theenumi}{\\Roman{enumi}} \\begin{enumerate} \\item That short GRBs originate from binary mergers of white dwarfs, neutron stars and/or black holes in all possible combinations \\cite{2006RPPh...69.2259M,2009ARA&A..47..567G}. \\item That long GRBs originate from the collapse of massive stars and the observed spikes in their light curves originate in the prolonged activity of an inner engine \\cite{1996ApJ...473..998F,1997ApJ...485..270S,1999ApJ...512..683F,1999PhR...314..575P}. In order to make the energy requirement less stringent they postulated the existence of a jet structure \\cite{1999ApJ...519L..17S,1999ApJ...526..707P}; in particular, the existence of a standard energetics for all GRBs of $\\sim 10^{51}$ erg \\cite{2001ApJ...562L..55F} was also claimed. \\item The observations of SNe associated only with long GRBs was considered a clear support for the idea that all long GRBs originate from supernovae in the collapse of very massive stars: the ``collapsar'' model \\cite{1993ApJ...405..273W,2006ARA&A..44..507W}. \\end{enumerate} \\renewcommand{\\theenumi}{\\arabic{enumi}}  Some of aspects of our work have offered a different point of view. We expressed them in three basic paradigms: \\begin{enumerate} \\item The relative space-time transformation (RSTT) paradigm \\cite{2001ApJ...555L.107R}: the first paradigm emphasizes the relevance of having the correct space-time parametrization of the source, which necessarily implies the knowledge of the equations of motion of the system and its entire worldline. This procedure was in contrast with the current practice of describing the GRB nature by a piecewise analysis. \\item The interpretation of the burst structure (IBS) paradigm \\cite{2001ApJ...555L.113R}: the second paradigm emphasizes the structure of the canonical GRB as composed of a proper GRB (P-GRB) and an extended afterglow, whose characteristics are mainly dominated by the two parameters of the total energy $E_{tot}$ and of the baryon loading $B=M_Bc^2/E_{tot}$ of the $e^+e^-$ plasma (where $M_B$ is the mass of the baryon loading): in the limit of $B \\to 0$ the canonical GRB would result in a short GRB and, in the opposite limit of $B \\to 10^{-2}$, which is the maximum possible value \\cite{2000A&A...359..855R}, it would result in a long GRB. Our model assumes spherical symmetry. \\item The GRB-supernova time sequence (GSTS) paradigm \\cite{2001ApJ...555L.117R}: the third paradigm introduces the concept of ``induced gravitational collapse,'' which considered the effect of a GRB triggering the late phase evolution of a highly evolved companion star leading to a supernova (SN). \\end{enumerate}  As we will see, the consistent application of these paradigms has led to conclusions quite different from the ones (I), (II), (III) formulated by the majority, especially with reference to short GRBs and to the collapsar model.  Before introducing the recent developments in our model, I outline some general considerations out of first principles establishing an upper limit on the black hole mass and, consequently, on the energetics of GRBs. I also recall a scenario outlined with John Wheeler. ",
        "conclusions": "the possibility that binary mergers can also lead to canonical long duration GRBs. It is clear, in fact, that if a binary system merges within a galaxy, then the average CBM density will be $\\sim 1$ particle/cm$^3$ and not the density typical of the halo (see Fig.~\\ref{picco_n=1}).     I have summarized the vast scenario made possible by the unprecedented technological possibilities of observing SNe, GRBs, GRBs associated with SNe, all the way to neo-neutron stars and proto-black holes. The corresponding theoretical work is rapidly expanding over new regimes of ultrarelativistic field theories and encompassing all fundamental interactions. A novel unified approach to nuclear physics, neutron stars, white dwarfs and gravitationally collapsed objects is rapidly emerging. All this offers unprecedented possibilities to enlarge our knowledge of fundamental physics and relativistic astrophysics, as well as to reach an understanding of yet unexplained phenomena of fundamental physics."
    },
    "1107/1107.2056_arXiv.txt": {
        "abstract": "{ In the last years, photometric and spectroscopic evidence has demonstrated that many, maybe all the Globular Clusters (GC) host multiple stellar populations. High-resolution spectroscopy has established that, while most GCs are mono-metallic with no significant abundance spread in $s$-elements, in all the globulars studied to date the presence of different stellar generation is inferred by the Na-O and the C-N anticorrelations.  In this context, NGC 1851 and NGC 6656 are among the most intriguing clusters. Contrary to the majority of GCs, they host two groups of stars with different s-elements abundance that are clearly associated to the two distinct sub-giant and red-giant branches detected in their color-magnitude diagrams (CMD). In the case of NGC 6656 s-rich stars are also enriched in iron and calcium. Each $s$-element group exhibits its own Na-O and C-N anticorrelations thus indicating the presence of sub-populations and suggesting that the parent clusters have experienced a very complex star-formation history. In this paper we summarize the properties of multiple populations in NGC 1851 and NGC 6656. ",
        "introduction": "In recent years, an increasing number of photometric and spectroscopic observational evidence have shattered the paradigm of globulars as the prototype of single, simple stellar populations (see Piotto \\ 2009 for a recent review).  Spectroscopic studies have demonstrated that most GCs are mono-metallic with no detectable spread in their iron content and also $s$-process elements do not exhibit large star-to-star variations in the majority of globulars  (e.g. Carretta et al.\\ 2009a, D'Orazi et  al.\\ 2010 and references therein). On the contrary, in all the clusters studied to date with large stellar samples, have been detected star-to-star variations in the light elements C, N, O, Na, and Al (e.g. Carretta et al.\\ 2009b, Pancino et  al.\\ 2010 and references therein). These variations are related to correlations and anticorrelations, which indicate the occurrence of high temperature hydrogen-burning processes which include CNO, NeNa, MgAl cycles and cannot occur in presently observed low mass GC stars.  Today it is widely accepted that the observed light-elements variations provide strong support to the presence of multiple stellar population in GCs with the second generations formed from the material polluted by a first generation of stars. On the contrary the debate on the nature of possible polluters is still open (e.g. D'Antona \\& Caloi al.\\ 2004, Decressin et al.\\ 2007).  While  abundance variations are well known since the early sixties, it was only the recent spectacular discovery of multiple sequences in the CMD of several GCs that provide an un-controversial prove of the presence of multiple stellar populations in GCs and brought new interest and excitement in GCs research (e.g. Piotto et al.\\ 2007). Photometric clues, often easy to detect simply by inspection of high-accuracy CMDs, arise in form of multiple main sequences (MS, Bedin et al.\\ 2004, Piotto et al.\\ 2007, Milone et al.\\ 2010), split sub-giant branch (SGB, Milone et al.\\ 2008, Anderson et al.\\ 2009, Piotto \\ 2009), and multiple red-giant branch (RGB, Marino et al.\\ 2008, Yong et al.\\ 2008, Lee et al.\\ 2009).  Apparently the main properties of stellar generations, like the relative fraction of stars, their location in the CMD, the spatial distribution, and the chemical behavior differ from cluster to cluster. % However, it is clear that some groups of clusters share all similar properties: \\begin{enumerate} \\item Na-O anticorrelation and C-N correlations are presents in all the GCs studied to date.  When stars with available Na and O abundances have been identified in the {\\it U} vs. ({\\it U-B}) CMD, it was found that the group of Na-poor (O-rich) stars are spread on the blue side of the RGB, while the Na-rich (O-poor) population define a narrow sequence on the red RGB (e.g. Marino et al.\\ 2008). \\item In few cases the MS morphology supports the presence of stellar populations with different helium. NGC 2808 has three distinct MSs (Piotto et al.\\ 2007) possibly associated to three stellar population with primordial helium and with Y$\\sim$0.33 and Y$\\sim$0.40 (D'Antona et al.\\ 2005). A spread MS have been detected in NGC 104 (47 Tuc, Anderson et al.\\ 2009) and  NGC 6752 (Milone et al.\\ 2010) where there is also some hint of a MS split. \\item In some GCs, like NGC 104, NGC 1851, NGC 5286, NGC 6388, and NGC 6656 there is a double SGB (Milone et al.\\ 2008,  Piotto \\ 2008,  Anderson et al.\\ 2009) associated to two stellar groups with either a difference in age by $\\sim$1 Gyr or with almost the same age but a significant difference in their overall C+N+O content (Cassisi et al.\\ 2008, Ventura et al.\\ 2009, Di Criscienzo et al.\\ 2010). \\item Finally there is the `extreme' case of $\\omega$ Centauri with either multiple MSs (e.g. Anderson\\ 1997, Bedin et al.\\ 2004), multiple SGBs (e.g. Sollima et al.\\ 2005, Villanova et al.\\ 2007), multiple RGBs (Lee et al.\\ 1999, Pancino et al.\\ 2000) and large star-to-star variation in iron and $s$-elements (e.g. Johnson \\& Pilachowski \\ 2010, Marino et al.\\ 2010).  Interestingly multiple populations have been detected also in the Sgr dwarf galaxy central cluster NGC 6715 (e. g. Siegel et al.\\ 2007, Piotto \\ 2009 and reference therein) and in Terzan 5  that is considered the surviving remnant of one of the primordial building blocks that are thought to merge and form galaxy bulges (Ferraro et al.\\ 2009).  \\end{enumerate}  NGC 1851 and NGC 6656 are among the most studied clusters of the third group and are the main subject of this paper. In the following we will compare the observed CMDs of these GCs and the abundance of the critical elements Fe, O, Na, and s-elements. ",
        "conclusions": "We have summarized some spectroscopic and photometric findings on multiple stellar populations in the GCs NGC 1851 and NGC 6656.  From spectroscopic and  photometric observations, it comes out that the SGB and RGB split is related to the presence of two groups of stars with different $s$-element contents and a possible difference in CNO abundance. According to this interpretation, $s$-poor stars would constitute the first population. The second stellar generation should have been formed after the AGB winds of this first stellar generation had polluted the protocluster interstellar medium with $s$-elements. In the case of NGC 6656, this second generation may have formed from material that was also enriched by core-collapse supernovae ejecta, as indicated by their higher iron, magnesium, and silicon content, and the lack of correlation of the iron content with a pure r-process element (Eu).  This scenario is seriously challenged by the recent findings on light elements Na, O, C, and N. As already mentioned in Sect.~1, the Na-O and the C-N anticorrelations can be considered the signature of multiple star-formation episodes. This implies that both clusters are composed by two groups of stars with distinct $s$-element content that define double SGB and RGB observed in the CMD but each group is made by multiple sub-populations of stars with different Na and O  abundance. NGC 1851 and NGC 6656 do not host only two stellar populations but have experienced a much more complex star-formation history with similarities to the `extreme' cases of $\\omega$ Centauri (Da Costa et al.\\ 2009, Da Costa \\& Marino, 2010) and the Sgr dwarf galaxy central cluster NGC 6715. A detailed comparison of the abundance patterns in  $\\omega$ Centauri and NGC 6715 with those of NGC 1851 and NGC 6656 is strongly needed to resolve this problem."
    },
    "1107/1107.2783_arXiv.txt": {
        "abstract": "{The recent detection of very high-energy (VHE) gamma-ray emission from the direction of the W43 star-forming region prompted us to investigate its stellar population in detail in an attempt to see wether or not it is possible an association.} {We search for the possible counterpart(s) of the gamma-ray source or any hints of them, such as non-thermal synchrotron emission as a tracer of relativistic particles often involved in plausible physical scenarios for VHE emission.} {We data-mined several archival databases with different degrees of success. The most significant results came from radio and near-infrared archival data.} {The previously known Wolf-Rayet star in the W43 central cluster and another cluster member appear to be resolved into two components,suggesting a likely binary nature. In addition,  extended radio emission with a clearly negative spectral index is detected in coincidence with the W43 cluster.  These findings could have important implications for possible gamma-ray emitting scenarios, which we also briefly discuss.} {} ",
        "introduction": "W43 (=G30.8$-$0.2)  is a star-forming complex first detected at radio wavelengths by \\cite{w58}. Recently, attention has been focused on it, because it is believed to be one of the closest starburst regions in our galaxy (see \\cite{b2010} and references therein). The distance to W43 was originally estimated based on H\\ion{I}~absorption and reddening measurements (\\cite{sbm1978, blum99}) pointing to a 5.5 kpc value, although with a significant uncertainty. An improved estimate using 6.7 GHz methanol masers yielded a more reliable average systemic velocity of 97.5 km $s^{-1}$, corresponding to a larger distance of $9.0 \\pm 0.4$ kpc (\\cite{pmg2008}).  A recent revision of survey molecular emission from W43 confirms a distance consistent with the 6 to 9 kpc range, with preference toward the lower limit, and a total mass of about $7\\times 10^6$ $M_{\\odot}$ (\\cite{nguyen}). We also refer the reader to this work for a wide-field view of the W43 as part of a very extended radio source in the galactic plane. From the cluster color-color diagram, \\cite{blum99} derived a considerable extinction value of $A_V \\simeq 34$ mag, equivalent to a hydrogen column density of about $6.5 \\times 10^{22}$ cm$^{-2}$.  The GLIMPSE image displayed in Fig. \\ref{tricromia+glimpse} shows the W43 field in the infrared band. Among other components, W43 contains a central open cluster of massive stars illuminating a giant H\\ion{II}~region, which will be better displayed below. A decade ago, \\cite{blum99}  studied the W43 cluster in detail showing that its three brightest members include one Wolf-Rayet (WR) and two luminous O-type stars. These objects are currently designated as  W43 \\#1, W43 \\#2 and W43 \\#3, respectively. The same authors classified the first of them as of WN7+abs spectral type and suggested that the presence of absorption features could be related to an unresolved companion. W43 \\#1 was subsequently  included in the 7th Catalogue of Galactic Wolf-Rayet stars (\\cite{vdh01}), where it is also listed as WR 121a.    \\begin{figure*} \\centering \\includegraphics[angle=0,width=15.0cm]{campograndeglimpse.eps} \\caption{Composite trichromatic image of the W43 cluster the GLIMPSE 3.6, 5.8 and 8 $\\mu$m bands in blue, green and red color, respectively. The green and yellow ellipses indicate the positions of the H.E.S.S. and Fermi LAT sources in the field. The W43 star-forming region, consistent with the H.E.S.S. position, is also indicated with a white label.} \\label{tricromia+glimpse} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[angle=0,width=15.0cm]{Fig2.eps} \\caption{Composite trichromatic image of the W43 cluster with the ESO VLT $J$, $H$ and $Ks$ bands in blue, green and red color, respectively. The positions of the W43 \\#1 and W43 \\#3 stars proposed binary components are indicated by arrows. White contours correspond to 6 cm radio emission as observed with the VLA in its extended B-configuration. Contours shown correspond to 4, 5, 6, 8, 10, 12, 15, 20, 25 and 30 times the rms noise of 0.76 mJy beam $^{-1}$. The VLA synthesized beam is plotted as an ellipse at the top right corner and corresponds to 1\\pri 63 $\\times$ 1\\pri 31, with position angle of 42\\grp 45.} \\label{tricromia+radio} \\end{figure*}  The W43 cluster is remarkably coincident in position with a VHE gamma-ray source detected by the H.E.S.S. collaboration at TeV energies (\\cite{chav2009}). This VHE object is known as \\hess\\ and it appears to be clearly extended with respect to the point-spread function of this Imaging Atmospheric Cherenkov Telescope array. If the association with the cluster is correct, this will be the second proposed case of gamma-ray emission from a stellar cluster after the well known Westerlund 2 (\\cite{aha2007}). W43 \\#1 (=WR 121a) has been also considered by the \\hess\\ discoverers as a possible counterpart candidate. Another gamma-ray source detected by the Fermi Large Area Telescope (LAT) is also present in the field (\\cite{0FGL,1FGL}).  However, its latest available LAT error box corresponding to the 1FGL J1848.1$-$0145c entry in the Fermi catalog excludes the W43 position at present. Whether the H.E.S.S. source is connected with the Fermi GeV emission or with the W43 stellar cluster still remains to be confirmed. Fermi has a strong background uncertainty in this region, which could imply significant position shifts in future improved versions of its catalog.  In this context, we report here new observational results about W43 mostly based on archival data. First, we present near-infrared observations that provide  strong evidence for the binary nature of the cluster WR star and another likely member as well. Secondly, we present the detection of extended non-thermal radio emission apparently coming from the cluster itself. Both could have strong implications in our understanding of possible gamma-ray emitting scenarios, and the third part of this paper is devoted to discuss this briefly. ",
        "conclusions": "The main results of this paper can be summarized as follows:  \\begin{enumerate}  \\item Based on archival data, we have reported the best angular resolution observations of the W43 central cluster at infrared and radio wavelengths, which shed some light on its possible association with high-energy gamma-ray emission.  \\item The binary nature of  the WR star W43 \\#1 has been confirmed by resolving its components with sub-arcsecond angular separation. One of them is clearly coincident with the  {\\it Chandra} X-ray source in the field. Another member of the cluster is also proposed to be a binary system consisting of two O-type stars.  \\item Non-thermal extended radio emission is clearly detected from the direction of the cluster. We speculate that it could be the result of a collective effect of the stellar winds from the whole population of luminous massive stars in the W43 cluster as the most conceivable scenario. Other alternative interpretations could also be considered and more observations are required to fully solve this issue.      \\end{enumerate}"
    },
    "1107/1107.5381_arXiv.txt": {
        "abstract": "We analyze the multi-wavelength photometric and spectroscopic data of 12 ultraluminous infrared galaxies (ULIRGs) at $z \\sim 1$ and compare them with models of stars and dust in order to study the extinction law and star formation in young infrared (IR) galaxies. Five extinction curves, namely, the Milky Way (MW), the pseudo MW which is MW-like without the 2175 \\AA\\ feature, the Calzetti, and two SN dust curves, are applied to the data, by combining with various dust distributions, namely, the uniform dust screen, the clumpy dust screen, the internal dust geometry, and the composite geometry with a combination of dust screen and internal dust. Employing a minimum  $\\chi^2$ method, we find that the foreground dust screen geometry, especially combined with the 8 - 40 $M_{\\sun}$ SN extinction curve, provides a good approximation to the real dust geometry, whereas internal dust is only significant in 2 galaxies. The SN extinction curves, which are flatter than the others, reproduce the data of 8(67\\%) galaxies better. Dust masses are estimated to be in excess of $\\sim 10^8 M_{\\sun}$. Inferred ages of the galaxies are very young, 8 of which range from 10 to 650 Myr. The SN-origin dust is the most plausible to account for the vast amount of dust masses and the flat slope of the observed extinction law. The inferred dust mass per SN ranges from 0.01 to 0.4 M$_{\\sun}$/SN. ",
        "introduction": "Recent developments in far-infrared(IR) and submillimeter astronomy have revealed new populations of galaxies associated with thermal dust emission, such as dusty quasars and infrared galaxies. Thermal dust emission detected in quasars and galaxies at high redshift exhibits inferred dust mass in excess of $\\sim 10^{8} M_{\\sun}$. Such a vast amount of dust observed at $z \\gtrsim 6$ (\\citealt{Bertoldi2003}; \\citealt{Robson2004};\\citealt{Beelen2006}) strongly suggests that the dust has been condensed in ejecta of core-collapsed supernovae (SNe II) promptly after the onset of the initial star formation. Asymptotic giant branch (AGB) stars are known to be the major source of dust in the present-day universe. However, the importance of AGB stars at high redshift is in controversy, because they inject dust grains into the ISM only after their progenitor stars have evolved off the main sequence. While a delay time of $\\gtrsim$ 500 Myr since their formation has been suggested (e.g., \\citealt{Dwek1998}), \\citet{Valiante2009} proposed that AGB stars can dominate dust production in a delay time of $\\gtrsim$ 150 Myr if the starburst is instantaneous. \\citet{Michalowski2010} pointed out that AGB stars are not efficient enough to form dust in dusty quasars at $z > 5$, whereas SNe II may be able to account for the dust, and underlined the importance of dust grain growth in the interstellar medium (ISM).  Quasars have been frequently used to study the extinction law at high redshift. This is because of their simple dust geometry, in which the central point-like emitting source is surrounded by dust grains which are physically separated from the source itself. This is analogous to the classical configuration assumed when dereddening single stars. There is a line of evidence that active galactic nuclei (AGNs) have extinction curves that differ from those observed in the Milky Way (MW), Large and Small Magellanic Clouds (LMC and SMC). \\citet{Gaskell2004} analyzed a large sample of AGNs and found indications of a very flat extinction curve in the ultra-violet(UV). On the other hand, exploring large samples of quasars from the Sloan Digital Sky Survey(SDSS), \\citet{Richards2003} and \\citet{Hopkins2004} found that the extinction law in quasars at $z < 2.2$ is described by SMC reddening but neither by LMC and SMC reddening nor by that found in the \\citet{Gaskell2004} sample. Recently, \\citet{Gallerani2010} proposed that quasars at $z = 3.9 - 6.4$ have reddening that deviates from that of the SMC with a tendency flatten at $\\lambda \\lesssim$ 2000 \\AA. \\citet{Maiolino2004} found that SDSS J1048+4647, a broad absorption line quasar (BALQSO) at $z$=6.2, requires a flat extinction curve which is quite different from those observed in the MW, LMC and SMC, but in excellent agreement with SN extinction curves by \\citet{Todini2001}.  Although it is reasonable to expect that the extinction law in young galaxies would be dominated by SN-condensed dust, the complications introduced by dust geometries have hampered this line of studies; the attenuation of extended starlight by dust depends not only on the optical properties of dust but also on the distribution geometry of the dust. \\citet{Calzetti2001} derived the empirical extinction law in local starbursts and blue compact galaxies, assuming the dust geometry described in a foreground screen, i.e., an apparent attenuation of starlight to be $e^{-\\tau_{\\nu}}$. If a dust geometry in their galaxies is indeed a foreground screen, the slope (i.e.,wavelength dependence) of the extinction curve (i.e., the optical depth $\\tau_{\\nu}$) is significantly flatter than those observed in the MW, LMC, and SMC. The SED and spectroscopic data of a young, ultraluminous infrared galaxy (ULIRG)\\footnote{ Infrared galaxies are classified according to bolometric infrared luminosity $L_{IR}(8 - 1,000\\micron)$; $L_{IR} > 10^{11} L_{\\sun}$ as luminous infrared galaxies (LIRGs), $L_{IR} > 10^{12}L_\\odot$ as ultraluminous infrared galaxies (ULIRGs),and $L_{IR} > 10^{13}L_\\odot$ as hyperluminous infrared galaxies (HyLIRGs) \\citep{Sanders1996}.} at $z \\sim 1$ were analysed by \\citet{Kawara2011}, who found that the data require a flat extinction curve and a dust geometry of a foreground dust screen, resulting in better agreement with the extinction curve synthesized from SN-condensed dust by \\citet{Hirashita2005} than those observed in local starbursts, MW, LMC, and SMC.  The slope of the extinction curve is expected to reflect the sites of dust formation and the subsequent processing in the ISM. The extinction curves can be ordered  as the SMC, LMC, MW, and Calzetti curves from the steepest to the flattest. The synthetic extinction curves of newly condensed dust grains in SN ejecta before the ISM reprocessing (hereafter called original SN extinction curves) are flatter than the Calzetti curve (\\citealt{Todini2001}; \\citealt{Nozawa2003}; \\citealt{Maiolino2004}; \\citealt{Hirashita2005}; see also Figure 3 in \\citealt{Kawara2011}). Then, the bulk of dust grains condensed in the ejecta is destroyed through the passage of the reverse shock (\\citealt{Bianchi2007};\\citealt{Nozawa2007};\\citealt{Hirashita2008}). Finally, grains that survive the reverses shock are shattered in the warm ionized ISM \\citep{Hirashita2010}. Nonetheless, the final SN extinction curves, after taking into account the reverse shock and shattering, are expected to have a slope as flat as those of the original extinction curves. The shattering may also be important in old galaxies like MW, LMC, and SMC, where AGB stars are dominant sources of dust production, because some observations (\\citealt{Groenewegen1997};\\citealt{Gauger1999}) suggest AGB stars form large grains. Dust grain growth in the ISM would be efficient in infrared galaxies where there must be many molecular clouds near SN remnants, and it is suggested that the dust mass in the high-redshift quasars is dominated by grain growth in the IMS (\\citealt{Dwek1998};\\citealt{Draine2009};\\citealt{Michalowski2010};\\citealt{Valiante2011}). It seems that this process result in flat extinction curves in the sense that the larger grains show the flatter extinction curve. However, the extinction laws found in the high-redshift quasars and local starbursts are flatter than those found in the MW for which \\citet{Draine2009} attributes the bulk of the dust mass to the growth of grains in the ISM.   As discussed so far, $z < 2.2$ quasars show a striking contrast in the slopes of the extinction curves against the $z \\sim 1$ ULIRG and higher redshift quasars. In this paper, we explore this issue by expanding our previous work on a single young ULIRG \\citep{Kawara2011} to 12 young ULIRGs at $z \\sim 1$. ",
        "conclusions": ""
    },
    "1107/1107.0271_arXiv.txt": {
        "abstract": "In this paper we extend the range of consistency of a constant bulk viscosity model to redshifts up to  $z\\sim 8.1$. In this model the dark sector of the cosmic substratum is a viscous fluid with pressure $p= -\\zeta \\theta$, where $\\theta$ is the fluid-expansion  scalar and $\\zeta$ is the coefficient of bulk viscosity. Using the sample of 59 high-redshift GRBs reported by Wei (2010), we calibrate GRBs at low redshifts with the Union 2 sample of SNe Ia, avoiding then the circularity problem. Testing the constant bulk viscosity model with GRBs we found the best fit for the viscosity parameter $\\tilde{\\zeta}$ in the range $0<\\tilde{\\zeta}<3$, being so consistent with previous probes; we also determined the deceleration parameter $q_0$ and the redshift of transition to accelerated expansion. Besides we present an updated analysis of the model with CMB5-year data and CMB7-year data, as well as with the baryon acoustic peak BAO. From the statistics with CMB it turns out that the model does not describe in a feasible way the far far epoch of recombination of the universe, but is in very good concordance for epochs as far as $z\\sim 8.1$ till present. ",
        "introduction": "Since 1998 reliable cosmological data have been accumulated leading to the conclusion that our universe has entered recently into an accelerated expansion epoch. The luminosity-distance data of supernovae Ia (SNe Ia) has provided the main evidence. Hitherto explanations to this phenomenon are not satisfactory. Most models consider accelerated expansion as due to a component of the universe that behaves opposite to gravity, the so called \\textit{dark energy} (DE). Another unknown component of the universe is the \\textit{dark matter} (DM), the missing mass necessary to held together galaxy clusters, also needed to explain the current large scale structure of the universe. The preferred model to describe accelerated expansion is the Lambda-Cold-Dark-Matter ($\\Lambda$CDM). However, cosmological constant has several theoretical drawbacks, like the discrepancies between the observed and theoretical values, i.e. the fine tuning problem; the coincidence problem is another objection, meaning the fact that dark energy is nearly equal to the matter density just now. Among many proposals to give a satisfactory answer there are some models that consider DE as a dynamical component, such as Quintessence or K-essence; in others proposals DE appears as the result of the interaction of fundamental particles. In here we adhere to the stream that considers DM and DE as different manifestations of the same component, describing the dark sector as some kind of matter whose physical properties depend on the scale: behaving as DM at high densities and transforming into DE at lower densities. These Unified Dark Matter models have the generalized Chaplygin gas as its paradigmatic example \\cite{Kamenshchik:2001, Bilic:2002, Bento:2002}, however these models present undesirable features like oscillations in the matter power spectrum and to mend this inconvenience it should be assumed the existence of nonadiabatic pressure perturbations.   On the other side, in the context of inflation of the very early Universe, it has been known since long time ago that an imperfect fluid with bulk viscosity can produce an acceleration in the expansion without falling back on a cosmological constant or some inflationary scalar field \\cite{Heller:1973, Heller:1975, Zimdahl:1996}. Inspired in those inflationary models there are some recent \\cite{Colistete:2007, Meng:2007, Ren:2007, Ren:2006, Ren:2006a, Meng:2009, Nucamendi:2009} and not so recent \\cite{Padma:1987} developments that assume a universe filled with a viscous single fluid. The bulk viscosity contributes to the cosmic pressure and drives the accelerated expansion.  The origin of the bulk viscosity in a physical system is due to its deviations from the local thermodynamic equilibrium \\cite{Israel:1963, Weinberg:1971} and different cooling rates \\cite{Zimdahl:1996a, Zimdahl:1997}. In a cosmological fluid, the bulk viscosity arises when the fluid expands (or contracts) too fast so that the system does not have enough time to restore the local thermodynamic equilibrium and then it appears an effective pressure restoring the system to its local thermodynamic equilibrium. Bulk viscosity can be seen as a measurement of this effective pressure. When the fluid reaches again the thermal equilibrium then the bulk viscosity pressure ceases \\cite{Wilson:2007}, \\cite{Okumura:2003, Ilg:1999, Xinzhong:2001}. So, in an accelerated expanding universe it may be natural to suppose that the expansion process is actually a collection of states out of thermal equilibrium in a small fraction of time that being modeled by a bulk viscous fluid might be a more realistic description of the accelerated universe today.  Moreover, these models have in their favour that behave well when density perturbations come into play: the power spectrum for viscous matter is well behaved and consistent with large scale structure data \\cite{Zimdahl:2010}. In particular, it does not suffer from the oscillation problem of Chaplygin models \\cite{Zimdahl:2009}.  On the other hand, since the earliest evidence of tight correlations in gamma-ray bursts spectral properties \\cite{Amati:2002}, the possibility arose of using GRBs as standard candles. Being so GRBs may open a window in redshift as far as $z \\sim 8$ \\cite{Salvaterra}, extending then the attainable range provided by SNe Ia observations. For an overview of most recent missions, see \\cite{McBreen} and references therein. At these epochs, $z \\sim 8$, the universe was dominated by dark matter, from which follows that this tool is less sensitive to dark energy. However for models where dark energy and matter are coupled \\cite{Amendola:2003, Bertolami:2007} or unified \\cite{Bento:2002} GRBs might be a useful tool.  GRBs were prevented of being used as standard candles because the intrinsic faintness of the nearby events, a fact that introduced a bias towards low redshifts of GRB and therefore the extrapolation of their correlations to low-z events faced serious problems. However this problem is solved if one uses SNe Ia data, in a combined calibration with GRBs at low redshifts, allowing then GRBs be considered as distance indicators \\cite{Ghirlanda, Amati:2008, Bertolami:2006, Liang:2006}. In particular, Liang \\cite{Liang:2008} implemented a method to circumvent this objection, obtaining the distance modulus of GRB at low redshift by interpolating from the Hubble diagram of SNe Ia. Then GRB relations are calibrated without assumming a particular cosmological model; fair is to say that there is some criticism to the method \\cite{Graziani}.  In this paper we test the model (presented in \\cite{Nucamendi:2009}) that considers a bulk viscous fluid as source of matter, with GRBs data given in \\cite{Wei:2010} containing redshifts up to $z=8.1$. In this way we extend the previous SNe Ia probes of the model. In Sec. II we introduced the main aspects of the bulk viscosity model; in Sec. III the calibration of GRBs is addressed. In Sec. IV we apply the calibrated relation to higher redshifts and do the statistics to find the best fit for the bulk viscosity parameter. We carry on the analysis with the observational data including SNe Ia; BAO for clusters with redshifts up approximately 2, as well as CMB5-year and CMB7-year data from WMAP; this last probe tests the model up to the last scattering surface,  about $z\\simeq1091.3$; according to our analysis the model is not reliable to such far epoch, but we extend its applicability to epochs of formation of structure $z \\sim 8$. We discuss conclusions and perspectives in the last section. ",
        "conclusions": "Modeling the dark sector of the universe with a bulk viscous fluid with pressure $p=- \\tilde{\\zeta} \\theta$, with a constant bulk viscosity $\\tilde \\zeta$, is tested with several cosmological data: SNe Ia, CMB and BAO jointly with the calibrated sample of GRBs. The model is in good agreement with observational data from SN Ia, BAO and GRBs, up to the range of $z=8.1$ with a chi-square statistics of $\\chi^2=1.01$.  Probing with SNe Ia and GRBs we obtained as the present Hubble parameter $H_0=69.56\\text{km}\\text{s}^{-1}\\text{Mpc}^{-1}$, the deceleration parameter $q_0=-0.47$, both in good agreement with the accepted values; as the redshift transition parameter we obtain $z_t=1.37$. When the baryon acoustic oscillations (BAO) data are included in the data-model confrontation, we obtain a viscosity parameter $\\tilde \\zeta=1.84$, included in the correct range, and improved values for $z_t=1.16$ and $\\chi^2=1.01$. The results with BAO confirm the analysis performed by Hip\\'olito-Ricaldi (2009) \\cite{Zimdahl:2009} on the matter power spectrum showing the compatibility of the bulk viscosity models with the 2dFGRS and SDSS data.  However the model does not pass the CMB test, exhibiting $\\chi^2=2.15$ for CMB5 and $\\chi^2=2.14$ for CMB7, besides a negative bulk viscosity constant, violating so the second law of thermodynamics. These results are consequence of neglecting a radiative component in the universe, that in so far epochs as the recombination, plays an important role. Thus, when testing so far epochs a radiation component should be added to the universe content, a component that is conserved independently of the rest of matter considered.  To conclude, constant bulk viscosity matter reproduce observational SN Ia, GRBs and BAO data in good agreement ($\\chi^2=1.01$), being then a reliable model not only for the recent state of the universe, but for epochs where structure already existed ($z=8.1$). It is an open questions if the model could account for the CMB data if a radiation component is included in the considered matter content of the universe."
    },
    "1107/1107.0792_arXiv.txt": {
        "abstract": "s{ We examine the neutralino dark matter (DM) phenomenology in supersymmetric scenarios with nonuniversal Higgs masses (NUHM) at the gauge coupling unification scale that can acommodate a light Higgs boson, where the correct relic density is obtained mostly through the annihilation into a pseudoscalar A. Our analysis shows that most part of the A pole region can produce detectable gamma-ray and antiproton signals. We further focus on uncertainties influencing the results in indirect and mainly direct detection.} ",
        "introduction": "\\subsection{The model}\\label{subsec:TheModel} One of the major experimental constraints on the constrained-MSSM parameter space comes from the LEP-2 limits on the lightest higgs boson mass. In particular, LEP-2 set a lower bound of $114.4$ GeV for this mass \\cite{hlim}, excluding the largest part of the model's viable parameter space. However, strictly speaking, this bound only applies to the SM. It comes mostly from searches in the Higgsstrahlung channel, which in the case of the MSSM is actually not identical to the SM one. In particular, $\\sigma_{\\mbox{\\begin{tiny}MSSM\\end{tiny}}}(e^+ e^- \\rightarrow hZ) = \\sin^2(\\beta - \\alpha) \\sigma_{\\mbox{\\begin{tiny}SM\\end{tiny}}}(e^+ e^- \\rightarrow hZ)$ \\cite{HaberNP,djouadihiggs}. Hence, the LEP2 bound $m_h \\gtrsim 114$ GeV applies to the MSSM only if $\\sin^2(\\beta - \\alpha) = {\\cal{O}}(1)$. This is the case in cMSSM/mSUGRA scenarios.  This limit can actually be partially circumvented once some of the cMSSM constraints are relaxed. In particular, it has been pointed out \\cite{NUHM1,NUHM2,mynusm} that relaxing the requirement for higgs mass universality at the GUT scale can effectively reduce the $\\sin^2(\\beta - \\alpha)$ factor, leading to a smaller cross-section and thus to weaker bounds on the lightest higgs mass. \\\\ Our particular model \\cite{Das:2010kb} is characterised by the following parameters \\begin{equation} m_{1/2}, \\ A_0,  \\ \\mbox{sign}(\\mu), \\ \\tan\\beta, \\ m_0, \\ m_{H_u}^2, \\  m_{H_d}^2 \\end{equation} where the GUT-scale common scalar mass $m_0$ concerns all scalars but the two Higgs bosons.  We examined the neutralino dark matter - related phenomenology of this model, notably the behavior of the relic density over some region of the parameter space, the prospects for indirect detection as well as the constraints coming from direct DM detection experiments. Viable parameter space points have to satisfy a number of constraints: \\\\ {{\\bf - Higgs boson mass limit:}} In the {\\it non decoupling} region where the $A$ boson becomes very light the lower limit of $m_h$ goes down to 93 GeV or even lower. We consider that the parameter space with $\\sin^2(\\beta-\\alpha) < 0.3$ (or, $\\sin(\\beta-\\alpha) \\lesssim 0.6$), and $93<m_h<114$ is in agreement with the LEP2 limit \\cite{hlim}. Consequently, the coupling of the heavier Higgs boson to the $Z$ boson ($g_{ZZH} \\propto \\cos(\\beta-\\alpha)$) becomes dominant and this makes the heavier Higgs boson SM - like, so the LEP-2 $114$ GeV  limit starts applying for the heavier $CP$-even higgs boson. On the other hand, in the decoupling region $\\sin(\\beta-\\alpha) \\sim 1$, which means that the $114$ GeV limit applies to the lightest higgs. Given the fact that there exists an uncertainty of about 3~GeV in computing the mass of the light Higgs boson \\cite{higgsuncertainty}, we accept a lower limit of $111$~GeV. \\\\ {\\bf - $Br(b\\rightarrow s\\gamma)$ constraint:} We demand \\cite{kim,bsg-recent} $2.77 \\times 10^{-4} < Br (b \\rightarrow s \\gamma) < 4.33 \\times 10^{-4}$. \\\\ {\\bf - $Br(B_s\\rightarrow \\mu^+ \\mu^-)$ constraint:} We further impose the important $Br(B_s \\to \\mu^+ \\mu^-)$ constraint coming from CDF, \\cite{CDF} ${\\rm Br} ( B_s \\to \\mu^+ \\mu^-) < 5.8 \\times 10^{-8}$ (at ${\\rm 95\\,\\%\\,C.L.}$), which has recently been improved to $< 4.3 \\times 10^{-8}$ at $95\\%$ C.L \\cite{Morello:2009wp}. \\\\ {\\bf - WMAP constraint :} In computing the relic density constraint, we consider the 3$\\sigma$ limit of the WMAP data \\cite{WMAPdata} $0.091 < \\Omega_{CDM}h^2 < 0.128$. Here $\\Omega_{CDM}h^2$ is the dark matter relic density in units of the critical density and $h=0.71\\pm0.026$ is the Hubble constant in units of $100 \\ \\rm Km \\ \\rm s^{-1}\\ \\rm Mpc^{-1}$. We use the code micrOMEGAS \\cite{micromegas} to compute the neutralino relic density.   \\subsection{Dark matter detection}\\label{subsec:DMdetection}  There exist two main modes of dark matter detection, usually referred to as ``indirect'' and ``direct''. Indirect detection is based on the principle that if DM can annihilate (or decay), in the early universe in order to give the measured relic density, this process should also occur today throughout the galaxy (and beyond), so we could hope to detect its annihilation products: gamma-rays, positrons, antiprotons and neutrinos. Direct detection of DM relies on the fact that WIMPs may interact with (scatter on) ordinary matter. This scattering is an in principle measurable effect and indeed a huge effort is currently being developped worldwide to measure potential signals coming from DM scatterings upon large underground detectors.  Concerning indirect detecion, in this work we compute the gamma-ray signals at intermediate galactic latitudes \\cite{Stoehr:2003hf} in the spirit of elliminating as much as possible uncertainties coming from the DM ``halo profile'' (i.e. its distribution in the galaxy) as well as background contributions to the spectrum. These gamma-rays can be detected by the Fermi satelite \\cite{Abdo:2010nz}. For the case of antiprotons, we compute the prospects for detection in the AMS-02 mission \\cite{Goy:2006pw}. We adopt a semi-analytical treatment of the diffusion equation \\cite{Maurin:2006hy,Maurin:2006ps,Lavalle:1900wn}, presenting results for the so-called MAX propagation model.  Finally, we compare the neutralino-nucleon spin-independent scattering cross-section to some of the tightest bounds available in the literature \\cite{cdms-2,xenon100,Kopp:2009qt}. As it has been pointed out, uncertainties are not absent in direct detection as well. These can be of a twofold nature: First of all, when giving exclusion bounds on the $(m_\\chi, \\sigma_{\\chi (p,n)})$ plane, a set of astrophysical assumptions (as well as some assumptions on the passage from the nuclear to the nucleonic level) have already been made. These assumptions are thought to have a small impact on the results, the mangitude of which depends, among other factors, on the considered mass range \\cite{McCabe:2010zh,Green:2010ri,Weber:2009pt,Salucci:2010qr,Vogelsberger:2008qb}. Secondly, there are often uncertainties in the cross-section computations performed by theorists in specific models. In our case, what is of relevance is one of the parameters entering the passage from the partonic to the the WIMP-nucleon scattering cross-section, denoted by $f_{T_s}$. This parameter is actually related to the strange quark content of the nucleon. Its value can be either measured or estimated through lattice QCD methods. The DarkSUSY \\cite{darksusy} code, which was used to compute the cross-section, adopts a default value of $f_{T_s} = 0.14$. However, recent lattice simulations \\cite {Cao:2010ph,lattice} point towards much lower values, of the order of $0.02$, being even compatible with zero. The effect of this uncertainty has been quantified \\cite{Ellis:2008hf} and is known to range from negligible to very large, depending on the specific mechanism driving the scattering cross-section. In what follows we shall quantify the effect of this uncertainty showing that it is really crucial in assessing the viability of our models. ",
        "conclusions": ""
    },
    "1107/1107.2935_arXiv.txt": {
        "abstract": "Non-ideal MHD effects play an important role in the gas dynamics in protoplanetary disks (PPDs). This paper addresses its influence on the magnetorotational instability (MRI) and angular momentum transport in PPDs using the most up-to-date results from numerical simulations. We perform chemistry calculations using a complex reaction network with standard prescriptions for X-ray and cosmic-ray ionizations. We first show that no matter grains are included or not, the recombination time is at least one order of magnitude less than the orbital time within $5$ disk scale heights, justifying the validity of local ionization equilibrium and strong coupling limit in PPDs. The full conductivity tensor at different disk radii and heights is evaluated, with the MRI active region determined by requiring that (1) the Ohmic Elsasser number $\\Lambda$ be greater than $1$; (2) the ratio of gas to magnetic pressure $\\beta$ be greater than $\\beta_{\\rm min}(Am)$ as identified in the recent study by Bai \\& Stone (2011), where $Am$ is the Elsasser number for ambipolar diffusion. With full flexibility as to the magnetic field strength, we provide a general framework for estimating the MRI-driven accretion rate $\\dot{M}$ and the magnetic field strength in the MRI-active layer. We find that the MRI-active layer always exists at any disk radius as long as the magnetic field in PPDs is sufficiently weak. However, the optimistically predicted $\\dot{M}$ in the inner disk ($r=1-10$ AU) appears insufficient to account for the observed range of accretion rate in PPDs (around $10^{-8}M_{\\odot}$ yr$^{-1}$) even in the grain-free calculation, and the presence of solar abundance sub-micron grains further reduces $\\dot{M}$ by one to two orders of magnitude. Moreover, we find that the predicted $\\dot{M}$ increases with radius in the inner disk where accretion is layered, which would lead to runaway mass accumulation if disk accretion is solely driven by the MRI. Our results suggest that stronger sources of ionization, and/or additional mechanisms such as magnetized wind are needed to explain the observed accretion rates in PPDs. In contrast, our predicted $\\dot{M}$ is on the order of $10^{-9}M_{\\odot}$ yr$^{-1}$ in the outer disk, consistent with the observed accretion rates in transitional disks. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5809_arXiv.txt": {
        "abstract": "\\noindent Microlensing surveys, which have discovered about a dozen extrasolar planets to date, have focused on the small minority of high-magnification lensing events, which have a high sensitivity to planet detection. In contrast, second-generation experiments, of the type that has recently begun, monitor continuously also the majority of low-magnification events. We  carry out a realistic numerical simulation of such experiments. We simulate scaled, solar-like, eight-planet systems, studying a variety of physical parameters (planet frequency, scaling of the snowline with stellar mass $R_{\\rm{snow}}\\propto M^s$), and folding in the various observational parameters (cadence, experiment duration), with sampling sequences and photometric error distributions taken from the real ongoing experiment. We quantify the dependence of detected planet yield on cadence and experiment duration, e.g., the yield is doubled when going from 3-hour to 15-minute baseline cadences, or from an 80-day-long to a 150-day-long experiment. There is a degeneracy between the snowline scaling index $s$ and the abundance of planetary systems that can be inferred from the experiment, in the context of our scaled solar-analog model. After 4 years, the ongoing second-generation experiment will discover of the order of 50 planets, and thus will be able to determine the frequency of snowline planet occurrence to a precision of 10-30\\%, assuming the fraction of stars hosting such planets is between 1/3 and 1/10, and a snowline index in the range $s=0.5$ to 2. If most planetary systems are solar analogs, over 65\\% of the detected planets will be \\textquoteleft Jupiters\\textquoteright, five in six of the detected anomalies will be due to a single planet, and one in six will reveal two planets in a single lensing event. ",
        "introduction": "\\label{sec:intro}  Among the techniques currently used to discover extrasolar planets, microlensing has some unique capabilities (see, e.g. \\citealt{Gould.2010.B}). Apart from the fact that microlensing probes for planets around stars over a large volume of the Galaxy, among both the bulge and the disk population, and can also discover free-floating planets (\\citealt{Sumi.2011.A}), it is sensitive to planets at projected host separations that are distinct from those probed by other methods, namely about 1 to 10 AU. This corresponds to the area beyond the \\textquotedblleft snowline\\textquotedblright\\ where gas and ice giants are likely to form. Microlensing is thus complementary to other methods, such as radial velocity or transit surveys, which are most sensitive to planets in close orbits around their host stars, and to direct imaging of planets, which probes larger separations (e.g. \\citealt{Crepp.2011.A}). To date, however, only about a dozen planets have been discovered via microlensing (\\citealt{Bond.2004.A}, \\citealt{Udalski.2005.A},  \\citealt{Beaulieu.2006.A}, \\citealt{Gould.2006.A}, \\citealt{Gaudi.2008.A}, \\citealt{Bennett.2008.A}, \\citealt{Dong.2009.A}, \\citealt{Sumi.2010.A}, \\citealt{Janczak.2010.A}, \\citealt{Miyake.2011.A}, \\citealt{Batista.2011.A}), out of the total $\\sim$550 known planets\\footnote{\\tt http://exoplanet.eu}, and the $\\sim$1200 candidates recently announced by the Kepler mission (\\citealt{Borucki.2011.A}).  The paucity of microlensing-discovered planets is a consequence of the surveying strategy, largely followed to date, which has focused on high-magnification lensing events. As summarized in \\citet{Gould.2010.B}, all currently known microlensing-discovered planets were first discovered by one or both of two wide-field, low-cadence microlensing surveys, OGLE (\\citealt{Udalski.2009.A}) and MOA (\\citealt{Sako.2008.A}). Events whose initial light curves suggested high peak magnifications triggered global networks of telescopes for photometric (and sometimes spectroscopic) follow-up of these specific events. Indeed, among the 12 known microlensing planets (11 systems, as OGLE-06-109L has two planets) eight were discovered in events with amplifications greater than 100, and three had medium magnifications (greater than 10). All of these events had long event time scales (40-60 days), whereas the typical event time scale is $\\sim$20 days (see Section \\ref{sec:sim_samp}, below), and high cadences were therefore not essential. Only one event, OGLE-05-169L (\\citealt{Beaulieu.2006.A}), was a low-magnification event, with an amplification of 3.0 and a relatively short event time scale of 11 days. The source, however, was a bright star with a baseline magnitude of $I=14.3$, which permitted triggering of the follow-up network nonetheless. All of the above events had peak $I$-band magnitudes brighter than 16.7.  Given the limited observational capabilities to date, the high-magnification follow-up strategy has been the optimal choice. Bright events can be followed up with high cadence using small telescopes, and these events are fortuitously also those with the highest sensitivity to planets. However, only 1\\% of all microlensing events have magnifications $>100$. Since the surveys of the Galactic bulge have discovered roughly 700 events per year in recent years, and apparently only a fraction of all stars host planets near their snowlines, only a couple of planets per year have been discovered, resulting in the current microlensing planet tally.  Microlensing survey statistics have been used in the past to constrain the frequency and properties of snowline-region planets by \\citet{Albrow.2001.A}, \\citet{Gaudi.2002.A}, and \\citet{Tsapras.2003.A}. More recently, \\citet{Sumi.2010.A} derived the mass-ratio distributions of cold exoplanets, and found that Neptune-mass planets are three times more common than Jupiters beyond the snowline. \\citet{Gould.2010.B} find from the high-magnification events that the frequency of Solar-like systems around lens stars is about 1/6. However, two main problems complicate these estimates. One is the small-number statistics they are based on. The other is the complex social process that determines which events are followed up, and how, a process which then needs to be modeled in order to estimate the intrinsic planet properties. \\citet{Gould.2010.B} argue that the very chaotic nature of the process results, in the end, in a randomization that minimizes selection effects and gives a high level of completeness. Nevertheless, a more controlled microlensing experiment would provide confidence that this is indeed the case.  It has been clear for some time that, to make progress and overcome these problems, a \\textquotedblleft second generation\\textquotedblright\\ (hereafter generation-II) microlensing survey (e.g., \\citealt{Gaudi.2009.A}) must be initiated, one in which a large fraction of all ongoing microlensing events toward the Galactic bulge -- both low and high magnification -- are monitored continuously with a cadence that is high enough to detect planetary perturbations in the light curves. Although only a fraction of low-magnification events will reveal planets, even if they are present around the lens star, the numbers of such events are much larger than high-magnification events. In the balance, such a survey, over several years, could yield a sample of planets that is competitive with those from other techniques. The requirements are straightforward -- a network of 1m-class telescopes situated so as to allow 24-hour coverage of the bulge, equipped with degree-scale imagers. In June-July 2010, we carried out a 6-week pilot version of such a generation-II experiment. Starting April 2011, we have begun a full generation-II experiment. In this paper, we perform numerical simulations of the possible outcomes of this and similar experiments.  We begin our study by considering a variety of physical possibilities for the abundance of planetary systems and their properties. We then ray trace through the systems to generate a library of microlensing light curves with the correct distribution of observed properties, such as peak magnification, peak magnitude, and timescale. The simulated light curves are then \\textquotedblleft observed\\textquotedblright\\ using the real observational parameters of our ongoing experiment, including sampling, weather gaps and photometric noise. Finally, the data are processed through a realistic detection pipeline, to recover the detection efficiency and the expected number of detections for each combination of physical and observational parameters. The results we present here can therefore serve to predict the power of an experiment to distinguish among models, and to optimize the experiment's yield. We note that, although generation-II experiments will discover many well-characterized planetary events, our focus is on the mere detection of planetary anomalies, even when they may end up not being unambiguously modeled.  Our methodology builds upon, and extends previous studies of the planet detection efficiencies of microlensing surveys (e.g., \\citealt{Gaudi.2000.A}; \\citealt{Bennett.2002.A}; \\citealt{Bennett.2004.A}; \\citealt{Yoo.2004.A}; \\citealt{Dominik.2010.A}; \\citealt{Penny.2011.A}). However, to our knowledge, ours is the first study to combine the following important elements in such a simulation: we consider full 8-planet solar-analog systems; we investigate the effect of the snowline radius of those systems, for a range of scaling dependences on mass; for our simulated light curve generation, we use the observed distributions of lensing event timescales, impact parameters, and peak magnitudes, and the distributions of lens mass, distance, and velocity implied by Galactic models; we use the real sampling patterns and photometric error distributions (including outliers) of the ongoing generation-II experiment that we simulate. Eventually, we will use the tools we develop here to constrain the actual physical parameter space of exoplanets probed by microlensing.  In Section \\ref{sec:gen_2_net}, we present the observational setup of our generation-II network. In Section \\ref{sec:sim_samp}, we describe the simulations, starting with the light curve generation process, and up to the detection procedure. In Section \\ref{sec:results} we discuss the results we obtain as a function of the various input parameters. A summary of our results is given in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this work, we have studied various aspects of generation-II microlensing experiments -- the optimization of the observational parameters of such experiments, and their sensitivity to planets for different physical assumptions about the planetary systems. Our main results are the following: \\begin{enumerate} \\item Long duration experiments with continuous high cadence are vital. Experiments with low cadences of 3 hours will detect only about half as many planets as experiments with cadences of 15 minutes, and likewise experiments lasting one-half, versus a full, bulge season. \\item The current approach of using a network of 1-m-class telescopes for a generation-II experiment is adequate for tracking the low-magnification events discovered by these same telescopes. We have quantified, for our particular network and its parameters, the decrease of the planet detection efficiency with the decrease of the peak brightness of an event. \\item The abundance of planetary systems is degenerate with the snowline scaling index. Conversely, with constraints on planetary fraction from other planet-hunting techniques, microlensing surveys could discriminate among different planet formation scenarios that predict different snowline scalings. \\item The majority of events with anomalies will be due to single planet deviations (83\\%), assuming Solar system distance and mass ratios, but the non-negligible remainder can reveal two planets. \\item For Solar-like systems, at least 65\\% of the detections will be of \\textquoteleft Jupiters\\textquoteright, due to both the mass and the location near the Einstein radius of this planet. \\item The current generation-II experiment which continuously monitors 8 $\\rm deg^2$ of the Galactic bulge with the highest microlensing event rate, and with a cadence of 15-30 minutes, will be able to determine, after four full seasons, the abundance of planets near the snowline to a precision of 10-30\\%. This, assuming the fraction of Galactic stars hosting such planets is between 1/3 to 1/10, and a snowline scaling power-law index in the range 0.5 to 2. \\item The main contaminant to planet detection, binaries, can largely be filtered out already at the detection stage, based on the longer deviation timescales that they produce. Additional characterization and modeling of individual events in the timescale overlap region between planets and binaries will further lower the contamination fraction. \\end{enumerate}  A full generation-II experiment, of the type we have simulated here, is already in progress, our results show that first estimates of the frequency and nature of snowline-region planets will be possible soon. To improve the analysis, our future simulations will include also higher-order dynamical effects (parallax, orbital motion, xallarap), a distributions of source sizes, limb darkening, and a wider range of physical planetary system possibilities. Together with results from other planet-search techniques, a clearer picture should emerge of the properties of planetary systems and the physics of their formation."
    },
    "1107/1107.5038_arXiv.txt": {
        "abstract": "Following the notion that the Titius--Bode rule (TBR) may also be applicable to some extra-solar planetary systems, although this number could be relatively small, it is applied to 55~Cancri, which is a G-type main-sequence star currently known to host five planets.  Following a concise computational process, we tentatively identify four new hypothetical planetary positions given as 0.081, 0.41, 1.51 and 2.95~AU from the star.  The likelihood that these positions are occupied by real existing planets is significantly enhanced for the positions of 1.51 and 2.95~AU in the view of previous simulations on planet formation and planetary orbital stability. For example, \\citet{ray08} argued that additional planets would be possible between 55~Cnc~f and 55~Cnc~d, which would include planets situated at 1.51 and 2.95 AU.  If two additional planets are assumed to exist between 55~Cnc~f and 55~Cnc~d, the deduced domains of stability would be given as 1.3--1.6 and 2.2--3.3~AU. The possible planet near 1.5~AU appears to be located at the outskirts of the stellar habitable zone, which is however notably affected by the stellar parameters as well as the adopted model of circumstellar habitability.  We also compute the distance of the next possible outer planet in the 55~Cnc system, which if existing is predicted to be located between 10.9 and 12.2~AU, which is consistent with orbital stability constraints.  The inherent statistical significance of the TBR is evaluated following the method by \\citet{lyn03}.  Yet it is up to future planetary search missions to verify or falsify the applicability of the TBR to the 55~Cnc system, and to attain information on additional planets, if existing. ",
        "introduction": "One of the more controversial aspects in the study of the Solar System as well as studies of selected extrasolar planetary systems is the application of the Titius--Bode rule (TBR).  Historically, the TBR was first derived and applied to the Solar System where it played a significant role in the search for new planetary objects (e.g., \\cite{nie72}). The discovery of Uranus by Herschel in 1781 and the largest object of the Mars--Jupiter asteroid belt, Ceres, by Piazzi in 1801 appeared to confirm the applicability and significance of the TBR.  The TBR historically given by the formula \\begin{equation} r_n = 0.4 + 0.3 \\times 2^n ,~~~~n = - \\infty, 0, 1, 2, 3, ... \\end{equation} (in AU) is, however, entirely inapplicable to Neptune and to the former Solar System planet Pluto\\footnote{Pluto is mentioned merely for historical reasons; it is often considered to highlight the limitations of the TBR for the Solar System.  Pluto is in the 3:2 mean-motion resonance with Neptune and is thus not a dynamically independent object.  Moreover, long-term integrations have uncovered that its orbit is chaotic yet stable over billion-year timescales \\citep{mal93}.}, which is still often considered to be an appropriate representative of the steadily increasing number of Kuiper Belt Objects (e.g., \\cite{jew01}).  A considerable amount of previous discussion revolves around the interpretation of the TBR for the Solar System, including its well-known insufficiency, see, e.g., \\citet{gra94}, \\citet{hay98}, \\citet{lyn03}, and \\citet{nes04} for previous studies largely based on statistical analyses.  For example, \\citet{lyn03} argued that it is not possible to conclude unequivocally that laws (or rules) of Titius--Bode type are, or are not, significant, a conclusion that also appears to be consistent with the work by \\citet{hay98}. Moreover, the main conclusion of \\citet{lyn03} is also shared by \\citet{nes04}.  \\citet{nes04} also pointed out the curiosity that if the Earth's distance is excluded from the sequence, then the probability of the occurrence of that sequence by chance is radically reduced from 95 to 100\\% down to 3\\%.  If no appropriate physical explanation for the TBR is identified, it is typically argued that the TBR may be merely a rule of chance (``numerology\"), fuelled by the psychological tendency to identify patterns where none exist \\citep{new94}.  On the other hand, there are previous detailed astrophysical and astrodynamical studies that potentially point to a physical background of the TBR, at least for the Solar System and (by implication) for a selected number of extrasolar planetary systems as well.  For example, \\citet{whi72} argued that jet streams may develop in a rotating gaseous disk at discrete orbital distances given by a geometric progression.  This result appears to be broadly consistent with the work by \\citet{sch84}, who pursued an analytic analysis of the structure, stability, and form of marginally stable axisymmetric perturbations of idealized rotating gas clouds.  His study showed that the radial parts of the expansion solutions obtained from the governing differential yield simple oscillating functions with TBR-type features.  A similar approach was undertaken by \\citet{gra94}, who argued that a Titius--Bode type law emerges automatically as a consequence of the scale invariance and rotational symmetry of the protoplanetary disc, and that such geometrical relationships are a generic characteristic of a broad range of physical systems.  \\citet{li95} investigated the nonlinear development and evolution of density disturbances in nebular disks on the basis of one of their previous studies.  Their analysis showed that the perturbed density is unstable with respect to self-modulation, leading to the formation of density field localization and collapse. In the case of a self-similar collapse, the perturbed density was found to increase with time and to form steady rings with very large but limited amplitudes.  The distances of these rings were found to follow a geometric progression akin to the TBR.  Also, in his summary about the final stages of planetesimal accumulation, \\citet{lis93} pointed out that ``[t]he self-limiting nature of runaway growth strongly implies that protoplanets form at regular intervals in semimajor axis.\"  This result appears to entail geometrical spacing of planets in planetary systems, potentially modified by their long-term orbital dynamical evolution, in possible agreement with (a modified version of) the TBR.  Another approach invoking statistical numerical experiments was adopted by \\citet{hay98}.  They chose to fit randomly selected artificial planetary systems to Titius--Bode type laws considering a distance rule inspired by the Hill stability of adjacent planets.  They did not identify a heightened significance of the TBR, except that its meaning is that stable planetary systems tend to be spaced in a regular manner. However, they indicate that their method could be used to identify potentially unstable planetary systems, especially if applied in conjunction with long-term orbit integrations. A further contribution to the interpretation of the TBR was made by \\citet{not96} and \\citet{not97} based on an approach of scale relativity and quantization of the Solar System.  It attempts to describe the Solar System in terms of fractal trajectories governed by a Schr\\\"odinger-like equation.  The physical background and justification of this approach, however, remain highly uncertain. A controversial but testable aspect of this theory is that it proposes the existence of one or two small planets between the Sun and Mercury, which to date have no support through  observations.  On the other hand, these studies point to the general possibility of ``empty\" orbits, which may also be of interest to generalized applications of the TBR as well.  In the following, we will apply a generalized version of the TBR, without subscribing to a specific physical interpretation, to the planetary system of 55~Cancri (= $\\rho^1$~Cnc = HD~75732 = HIP~43587 = HR~3522; G8~V) (see Table~1), which was historically the fourth star other than the Sun after 51~Peg, 70~Vir, and 47~UMa (not counting the pulsar PSR~1257+12 as well as other controversial cases) that was identified hosting a planet.  A previous effort to apply the TBR to 55~Cnc has been pursued by \\citet{pov08}, which will also be discussed in the following.  Additionally, we will compare our results to a modified version of the TBR given by \\citet{lyn03} previously applied to the Solar System.  55~Cnc is known to host five planets, named 55~Cnc~b to 55~Cnc~f, which were discovered between 1996 and 2007 by \\citet{but97}, \\citet{mar02}, \\citet{mca04}, and \\citet{fis08}; note that all planets have been detected using the radial velocity method.  The positions of the planets, according to their original determinations, are given as 0.038, 0.115, 0.240, 0.781, and 5.77~AU, respectively, and their masses ($M_p \\sin i$) range from 0.034 to 3.835~$M_{\\rm J}$ (see Table~2; it also gives information on the respective uncertainties).  However, a notable controversy emerged about the innermost planet 55~Cnc~e.  \\citet{daw10} argued that the position of 55~Cnc~e is given as 0.016 rather than 0.038~AU based on an improved method that allowed to distinguish an alias from the true planetary frequency in a more reliable manner.  This revised distance was subsequently confirmed by \\citet{win11} based on continuous photometric monitoring with the MOST (Microvariability \\& Oscillations of STars) space telescope.  Moreover, there is an ongoing discussion about the principal possibility of additional planets in the 55~Cnc system, particularly in the gap between 0.8 and 5~AU.  This speculation is fuelled by detailed orbital stability studies, including studies involving putative Earth-mass planets \\citep{blo03,riv07,ray08,smi09,ji09}.  If such a planet exists, it would also possibly imply the presence of a potentially habitable planet if appropriate conditions are met (e.g., \\cite{lam09}).  The star 55~Cnc is a middle-aged main-sequence star with a mass of approximately 0.95 $M_{\\solar}$ \\citep{val05}, with a stellar effective temperature of about 5250~K (see Table~3).  From the {\\it Hipparcos} parallax of $79.8 \\pm 0.84$~mas \\citep{esa97}, a luminosity of $0.61 \\pm 0.04$ $L_{\\solar}$ (see Table~1) can be deduced, which is consistent with its spectral type G8~V \\citep{gonz98}; see also discussion by \\citet{mar02}.  The age of the star can be obtained from the strength of the chromospheric Ca~II H+K emission, indicating an age of $4.5 \\pm 1$~Gyr \\citep{don98,hen00}; see also \\citet{bal97} for previous work.  Another element of our work is the assessment of 55~Cnc's circumstellar habitability, which will be pursued following the approach by \\citet{kas93} and \\citet{und03}; additionally, it will consider the possible extension of the outer boundary of habitability as discussed by, e.g., \\citet{for97}, \\citet{mis00}, and \\citet{hal09}.  Our paper is structured as follows.  After the introduction of historical aspects of the TBR and 55~Cnc itself, as done, we describe our methods and results.  Emphasis will be placed on the calculational process concerning the TBR based on four different measures of difference; they are applied to assess the deviations between the positions of the five known planets of 55~Cnc (i.e., 55~Cnc~b to 55~Cnc~f) and the predictions of the TBR, thus allowing to constrain the unknown TBR parameters.  In addition, we comment on the possibility of a habitable planet in the 55~Cnc system as implied by the TBR.  Thereafter, we consider other relevant studies, particularly studies regarding constraints from orbital stability simulations and planet formation. We also comment on a previous application of the TBR to the 55~Cnc system given by \\citet{pov08}.  Moreover, we provide an analysis of the statistical significance for the TBR forwarded by this work, as well as of the previous TBR version by \\citet{pov08}, following the method of \\citet{lyn03}. Finally, we present our summary and conclusions. ",
        "conclusions": "The goal of our study is to provide an application of the TBR to the 55 Cancri star--planet system.  It is clearly understood that the TBR does not have the same standing as the various well-established methods for the search and identification of extrasolar planets; see, e.g., \\citet{jon08} for a detailed overview.  However, as pointed out in the Introduction, the TBR appears to be successful (within limits) for a large segment of the Solar System.  Based on principal considerations, it would therefore be highly extremely unlikely if there was no other planetary system for which a TBR-type planet spacing would be realized, although the number of those systems may be relatively small.  Thus, the predictions conveyed in this study follow the hypothesis that the TBR is applicable to the 55~Cnc planetary system.  This approach entails the prediction of four planetary positions interior to the outermost planet 55~Cnc~d, which are given as 0.081, 0.41, 1.51, and 2.95~AU.  Whether these positions are ``empty\", filled with asteroid / comet belts (like in one instant of the Solar System, where Ceres is customarily viewed as substitute for a planet) or occupied by a developed planet is beyond the scope of this work; in this case additional astrophysical requirements need to be met.  Pertaining to this latter point, the strongest case for additional planets around 55~Cnc is given for the possible planets near 1.51 and 2.95~AU, noting that they would be most compatible with detailed studies of orbital stability \\citep{ray08,ji09} as well as previous studies about 55~Cnc's proposed protoplanetary disk \\citep{fis08}.  For example, \\citet{ray08} argued that additional planets could exist between 55~Cnc~f and 55~Cnc~d, including planets located near 1.51 and 2.95 AU in the view of long-term orbital stability simulations.  Specifically, if two additional planets are assumed to exist between 55~Cnc~f and 55~Cnc~d, the domains of stability were derived as 1.3--1.6 and 2.2--3.3~AU.  These two ranges are fully compatible with the results of the present study.  Moreover, the possible planet near 1.51~AU is of particuar interest as it appears to be located at the outskirts of the stellar HZ.  The existence of this planet is consistent with both orbital stability simulations and studies of the formation of terrestrial planets for 55~Cnc.  For example, calculations of terrestrial planet formation by \\citet{ray06} showed that the majority of their simulations lead to the build-up of terrestrial planets in 55~Cnc's HZ, including planets near 1.5~AU from the star.  Nonetheless, the final verdict of potential habitability of this possible planet (if existing) depends on a sizeable number of factors, including the definition of the zone of circumstellar habitability, considering that the outer boundary of the stellar HZ is notably impacted by a variety of astrobiological aspects (e.g., \\cite{lam09}) including the structure, density and compositions of the planetary atmosphere (e.g., \\cite{for97,mis00,hal09}).  The TBR as implemented follows closely the previous version as applied to the Solar System with $A$, $B$, and $Z$ as free parameters, which were obtained through careful statistical fitting considering the five known planets of the 55~Cnc system.  Here we obtain $A \\simeq 0.031$, $B \\simeq 0.049$, and $Z \\simeq 1.98$. The parameter $B$ indicates the compactness of the planetary system, which is much higher than for the Solar System as also indicated by the conventional version of Solar System's TBR ($B=0.3$). The parameter $Z$ indicates the relative geometrical spacing of the planetary system; it is virtually identical to that of the Solar System ($Z=2.0$). We also computed the orbital distance of a the next possible outer planet of the 55~Cnc system, i.e., exterior to 55~Cnc~d. Our study predicts that if existing it would be located between 10.9 and 12.2~AU from the star.  This distance is consistent with previous orbital stability analyses, which predict the position of a possible exterior planet as beyond 10~AU \\citep{ray08}.  In summary, it is obviously up to future observational efforts to verify or falsify the existence of any possible additional planets. From a general perspective, these efforts will assist in gauging the applicability of the TBR outside of the Solar System, and will help to discriminate between different versions of the TBR if found applicable to the 55~Cnc star--planet system.  \\bigskip  The author acknowledges valuable comments by an anonymous referee, which also resulted in the addition of Section~4."
    },
    "1107/1107.1936_arXiv.txt": {
        "abstract": "The bright star 55~Cancri is known to host five planets, including a transiting super-Earth. We use the CHARA Array to directly determine the following of 55~Cnc's stellar astrophysical parameters: $R=0.943 \\pm 0.010 R_{\\odot}$, $T_{\\rm EFF} = 5196 \\pm 24$ K. Planet 55~Cnc f ($M \\sin i = 0.155 M_{Jupiter}$) spends the majority of the duration of its elliptical orbit in the circumstellar habitable zone (0.67--1.32 AU) where, with moderate greenhouse heating, it could harbor liquid water. Our determination of 55~Cancri's stellar radius allows for a model-independent calculation of the physical diameter of the transiting super-Earth 55~Cnc e ($\\simeq 2.1 R_{\\earth}$), which, depending on the assumed literature value of planetary mass, implies a bulk density of 0.76 $\\rho_{\\earth}$ or 1.07 $\\rho_{\\earth}$. ",
        "introduction": "55 Cancri (= HD~75732) is a late G / early K dwarf / subgiant currently known to host five extrasolar planets with periods between around 0.7 days and 14 years and minimum masses between 0.026 and 3.84 $M_{Jupiter}$. The super-Earth 55~Cnc e was recently discovered to transit (Winn et al. 2011, Demory et al. 2011), prompting a number of studies of the properties of this system (e.g., Kane et al. 2011, von Braun et al. 2011b). We used the CHARA Interferometric Array to directly measure the stellar angular diameter, which, when combined with Hipparcos parallax measurement and calculation of bolometric flux based on spectral templates and literature broad-band photometry, yields the physical $R_{star}$ and $T_{\\rm EFF}$ (von Braun et al. 2011a). ",
        "conclusions": ""
    },
    "1107/1107.3044_arXiv.txt": {
        "abstract": "Based on the database with 56 supernovae (SNe) events discovered in 3838 galaxies of the southern hemisphere, we compute the rate of SNe of different types along the Hubble sequence normalized to the optical and near-infrared (NIR) luminosities as well as to the stellar mass of the galaxies. We find that the rates of Type Ia SNe show a dependence on both morphology and colors of the galaxies, and therefore, on the star-formation activity. The rate of SNe~Ia can be explained by assuming that at least 15\\% of Ia events in spiral galaxies originate in relatively young stellar populations. We also find that the rates show no modulation with nuclear activity or environment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.3758_arXiv.txt": {
        "abstract": " ",
        "introduction": "} The Galactic OB-associations offer a unique opportunity to study the influence of massive stars on the interstellar matter. A reconstruction of the star-formation history of many Galactic fields should be possible once the spatial distribution of the young stars is reliably determined. Despite the extensive efforts to improve and unify the distances to the young stellar groups in the Milky Way (MW), discrepancies still remain in the published studies for a large number of fields. Many of the present distance estimates are based to a large extent on preliminary distance calibrations, broad-band photometry, or absolute magnitudes ($M_V$) obtained via spectral and luminosity type (MK classification). On the other hand, the $uvby\\beta$ photometric system provides  $M_V$ and colour excess determinations for early-type stars in excellent agreement with the $Hipparcos$ parallaxes \\cite{kaltcheva98}. \\begin{figure} \\center \\includegraphics[scale=0.27]{Kaltcheva_N_Fig1a.pdf} ~ \\includegraphics[scale=0.27]{Kaltcheva_N_Fig1b.pdf} \\caption{\\label{fig1} Comparisons of  $uvby\\beta$ $M_V$ (left) and MK-based $M_V$ (right) to the $Hipparcos$-based $M_V$.} \\end{figure}  Figure~\\ref{fig1} represents the comparison of $uvby\\beta$ $M_V$ and MK-based $M_V$ to the $Hipparcos$-based $M_V$, pointing out to a possible over-estimation of stellar distance when relying on a MK-based determination. Since the distances to the Galactic OB associations are based mainly on individual stellar distances and rarely on main-sequence fitting, applying $uvby\\beta$ photometry should lead to a significant improvement for many OB groups in the MW.  Our deriving of $uvby\\beta$ photometric distances utilizes the intrinsic colour calibrations of Crawford \\cite{crawford78} and Kilkenny \\& Whittet \\cite{kilkenny85} and the luminosity calibration of Balona \\& Shobbrook \\cite{balona84} and takes into account possible stellar emission and mis-classification \\cite{kaltcheva00}. The expected errors for one star are about 12 \\% for luminosity classes III-IV and about 18-20~\\% for luminosity classes I and II. Photometric $uvby\\beta$ distances derived in this way provide the same impression of the star-forming field's structure as the improved $Hipparcos$ parallaxes (cf. \\cite{kaltcheva07}).  In this contribution we present improved distance estimates for four Galactic OB associations. ",
        "conclusions": "As already mentioned, a better understanding of the phenomena associated with the interaction between young massive stars and their surrounding ISM in the fields of the Galactic OB associations should be possible once their structure is reliably established. A multi-wavelength approach analyzing the distribution of the massive OB stars, ionized and neutral material, and that of the interstellar dust would allow us to better understand the different components of the ISM and the interactions among them.  \\small  %"
    },
    "1107/1107.3272_arXiv.txt": {
        "abstract": "We describe  a dedicated cosmic-ray telescope that explores a new method for detecting Cerenkov radiation from high-energy primary cosmic rays and the large particle air shower they induce upon entering the atmosphere. Using a  camera comprising 16 multi-anode photomultiplier tubes for a total of 256 pixels, the Track Imaging Cerenkov Experiment (TrICE) resolves substructures in particle air showers with $0.086^{\\circ}$ resolution. Cerenkov radiation is imaged using a novel two-part optical system in which a Fresnel lens provides a wide-field optical trigger and a mirror system collects delayed light with four times the magnification. TrICE records well-resolved cosmic-ray air showers at rates ranging between 0.01-0.1 Hz. ",
        "introduction": "High-precision cosmic-ray composition measurements carried out by balloon- and satellite-borne experiments can achieve charge resolutions of less than 5\\% \\citep{2009ApJ...707..593A,2008ApJ...678..262A} and in some cases, can measure the isotopic abundance of these particles \\citep{0004-637X-698-2-1666}. Such experiments become limited by statistics at the few TeV/amu regime, because of their small geometric factors. On the other hand, ground-based arrays can measure the composition of cosmic rays from the few TeV through EeV regime by employing the Earth's atmosphere as a calorimeter and measuring properties of the extensive air shower (EAS) produced when a high-energy cosmic ray interacts with nuclei in the atmosphere \\cite{2009NuPhS.190..240T,Antoni20051}. Such indirect composition measurements rely on hadronic interaction models and have limited charge resolution. A new technique proposed by  \\citep{Kieda:2000ky} and implemented by \\citep{2007PhRvD..75d2004A,2010PhDT........37W} shows that by measuring the Cerenkov light initiated both directly by the cosmic ray (DC light) and by the secondary particles in the air shower (the EAS light), the composition of cosmic rays at TeV-PeV energies can be measured with $<20\\%$ charge and energy resolution.  Note that while the field of view of typical satellite- and balloon-borne experiments is typically greater than imaging Cerenkov telescopes, the ground-based arrays can also expand their exposures by collecting data for longer periods of time.  \\begin{figure}[htbp] \\centering \\includegraphics[width=90mm]{fig01.pdf} \\caption[The timing and angular separation between DC light and EAS light]{Simulated photons emitted via Cerenkov radiation starting from the top of the atmosphere collected in a ring between 50 m and 115 m from the shower impact.. The DC component is cleanly separated from the EAS light when the following quantities are considered: the angle between the cosmic ray's trajectory and the arrival direction of the photon and the time delay from when the cosmic ray would have impacted the ground and the arrival time of the photon on the ground. The DC component is compact in both time and angular space ($0.05^{\\circ}$ in angular extent and 1 ns in duration) and cleanly separated from the extensive air shower component by 2 ns and $0.2^{\\circ}$ in a simulated 50 TeV $^{56}$Fe air shower that is vertically incident and interacts 25 km above ground-level.  } \\label{fig:dc_timeangle} \\end{figure}  High-energy cosmic rays continuously produce Cerenkov radiation in the atmosphere until they interact hadronically with nuclei in the atmosphere. Thereafter, the cosmic ray generates an extensive air shower (EAS) producing $>10^{5}$ particles per shower, largely depending on the primary's energy. The light directly produced by a heavy nucleus is appreciable ($ > 3 \\times 10^{4}$ for the simulation in Fig. \\ref{fig:dc_timeangle}), but the difficulty lies in separating it from that produced in the extensive air shower. Above the first interaction height, the particle generates DC light with a narrow Cerenkov angle ($\\sim0.1^{\\circ}$), because the density--and therefore, the index of refraction which defines the Cerenkov angle--is low at the top of the atmosphere. Comparatively, light emitted in the extensive air shower will be emitted at broader angles with respect to the particle trajectory ($\\sim0.2$--$2.0^{\\circ}$). The DC light will also arrive $\\sim$2 ns after the extensive air shower, because it has to traverse more material than the photons produced lower in the atmosphere. The simulation illustrated by Fig. \\ref{fig:dc_timeangle} shows that by exploiting these two characteristics, the DC light can be identified.  The simulations were done using the Corsika Monte Carlo package, 6.2040 at an observation level of 1275 m above sea level \\citep{1998FZKA6097}.  Cleanly separating the DC light from the EAS light is critical, because the DC light can be used to measure the charge of the incoming particle while the EAS light can be used to measure initial energy of the cosmic ray. The number of Cerenkov photons produced by an individual charged particle scales with the square of the charge of the particle \\cite{0508-3443-6-7-301}.  Thus the measurement of the cosmic ray's electric charge depends only weakly on the hadronic interaction model used to simulate the events.  This technique was recently confirmed with the first ground-based detection DC light and the subsequent measurement of the energy spectrum of  five charge bands from protons to iron \\citep{2007PhRvD..75d2004A}.  The DC signal is compact in both time and angular space, while the extensive air shower is more broad. The variation in arrival times and angles in the EAS light arise because the air shower light is emitted at a wide range of heights and from a number of secondary particles. A telescope that makes use of the timing and angular separation between the two signals will be able to measure the species of the primary particle to high energy, provided that it has enough effective area to account for the diminishing flux at higher energies.  \\begin{figure}[htbp] \\centering \\includegraphics[width=90mm]{fig02.pdf} \\caption[The TrICE Telescope]{The completed TrICE telescope with the baffle and electronics rack in the center of the telescope. Spherical mirrors with 4-m focal length line the base of the structure supporting the camera, Fresnel lens, and planar mirror. Each spherical mirror has a 1-m diameter. The height of the telescope, from the baseplate to the top of the planar mirror is 3 m.} \\label{fig:trice_fulldet} \\end{figure}  The Track Imaging Cerenkov Experiment (TrICE), sited at Argonne National Laboratory and shown in Fig. \\ref{fig:trice_fulldet}, serves as a test bed for new technologies that exploit the timing and angular characteristics of DC events. In particular, it employs a novel optical trigger and a camera composed of multi-anode photomultiplier tubes (MAPMTs) to image cosmic-ray events in both high- and low-resolution modes. The primary objectives were to verify the proposed technique via the detection of the Cerenkov light of air showers and to engineering a test facility for camera technology, while the secondary objective was to observe DC events.  \\begin{figure}[htbp]    \\centering  \\subfloat[][Photons collected in an IACT-like Camera]{\\label{fig:IACTcamera}\\includegraphics[width=2.5in]{fig03a.pdf}}\\\\ \\subfloat[][Photons collected in a TrICE-like Camera]{\\label{fig:TrICEcamera}\\includegraphics[width=2.5in]{fig03b.pdf}} \\caption[Imaging DC light with TrICE and 2$^{nd}$-generation IACTs]{A comparison of the ability for a telescope with angular resolution comparable to a typical IACT camera (0.15$^{\\circ}$, (a)) and a telescope with a TrICE-like camera (0.086$^{\\circ}$, (b)) to separate DC light from EAS light from the simulation shown in Fig. \\ref{fig:dc_timeangle}, landing 70 m from the telescope. The DC light in the front of the shower--indicated here by the $+$ and $\\times$ symbols--is more prominent in the TrICE-like camera, because of the improved angular resolution. The pixel marked with the $+$ sign in (a) is composed of 57\\% DC light; the pixel marked with a $\\times$ sign, 10\\%. Similarly, the pixel marked with the $+$ sign in (b) is composed of 89\\% DC light; the photons in the pixel marked with a $\\times$ sign, 33\\%.}      \\label{fig:em_ang_comp}  \\end{figure} Imaging air Cerenkov telescopes (IACTs) are designed to map the longitudinal development of an EAS into angular space \\citep{2008RPPh...71i6901A}.  As demonstrated by the discovery of DC radiation by the H.E.S.S. collaboration, DC light can be detected by IACTs with 0.15$^{\\circ}$  angular resolution \\cite{2007PhRvD..75d2004A}. However, improving the angular resolution reduces the contamination of secondary air shower photons into the pixels containing mostly DC light, and therefore, reduces the dependence of the charge measurement on energy. To compare the power of the TrICE camera to a standard IACT, consider a  vertically-incident iron nucleus of energy 50 TeV that interacts 25 km above the observation level (Fig. \\ref{fig:em_ang_comp}). When the IACT images it, the DC signal is easily visible in the camera plane and concentrated into one phototube, towards the front of the shower. However, the timing separation ($\\sim$2 ns) can be ambiguous in the case of non-isochronous mirrors, and while the DC signal is visible, it is generally contaminated by air shower light because the angular separation between the DC light and the end of the shower, 0.2$^{\\circ}$, is only slightly larger than the 0.15$^{\\circ}$ viewing angle of a standard IACT photomultiplier tube \\citep{2008ICRC....2..417W}. The TrICE telescope, with its improved angular resolution of 0.086$^{\\circ}$, however, would image the DC light in multiple pixels and enhance the time separation between the two signals using the optics shown in Fig. \\ref{fig:trice_optics_demo}. ",
        "conclusions": "\\label{sec:trice_outlook} The main successes of the TrICE telescope include the stable operation of a MAPMT camera in a telescope employing the imaging atmospheric Cerenkov technique and the demonstration of a novel optical system, including a proof-of-concept automated mirror alignment system. The technologies tested during the construction and operation of this telescope may be considered for use in the design of an observatory dedicated to cosmic-ray observations.  The optical system provided both an optical trigger and the convenience of a fixed-mount design. Alignment of the fixed-mount optical system proves to be difficult when the telescope must be operated in weather conditions, but using a remotely-controlled alignment system and a fixed point-source, relative alignment of the mirrors has been shown to produce a point-spread function smaller than the MAPMT pixel size.  Note that because of the improved angular resolution, the optical quality must be better than what is found in typical IACTs.  The TrICE telescope demonstrates that a MAPMT camera can be stably operated in a cosmic-ray telescope under high ambient light conditions with little variation in operating conditions.  Well-resolved images of the cosmic ray air showers were seen in the TrICE data, demonstrating the level of precision available to the next generation of $\\gamma$-ray and cosmic-ray telescopes that will use cameras with comparable angular resolution.  Although no DC events were detected in TrICE, this is consistent with the limited $\\sim200$ m$^{2}$ s sr exposure which was obtained with the instrument. Upon the conclusion of the 2007 observing season, efforts were shifted from trying to identify DC events with TrICE to using the DC technique to measure the iron spectrum in VERITAS data \\citep{2010PhDT........37W}. Nevertheless, the lessons learned from the operation of this instrument can help guide the design of a future DC instrument, be it a dedicated observatory \\citep{2008ICRC....2..349W} or a the planned next-generation IACT array, the Cherenkov Telescope Array (CTA) \\citep{2010arXiv1008.3703C}."
    },
    "1107/1107.2454_arXiv.txt": {
        "abstract": "Our current understanding of radio-loud AGN comes predominantly from studies at frequencies of 5\\,GHz and below. With the recent completion of the Australia Telescope 20\\,GHz (AT20G) survey, we can now gain insight into the high-frequency radio properties of AGN. This paper presents supplementary information on the AT20G sources in the form of optical counterparts and redshifts. Optical counterparts were identified using the SuperCOSMOS database and redshifts were found from either the 6dF Galaxy survey or the literature. We also report 144 new redshifts. For AT20G sources outside the Galactic plane, 78.5\\% have optical identifications and 30.9\\% have redshift information. The optical identification rate also increases with increasing flux density. Targets which had optical spectra available were examined to obtain a spectral classification.  There appear to be two distinct AT20G populations; the high luminosity quasars that are generally associated with point-source optical counterparts and exhibit strong emission lines in the optical spectrum, and the lower luminosity radio galaxies that are generally associated with passive galaxies in both the optical images and spectroscopic properties. It is suggested that these different populations can be associated with different accretion modes (cold-mode or hot-mode). We find that the cold-mode sources have a steeper spectral index and produce more luminous radio lobes, but generally reside in smaller host galaxies than their hot-mode counterparts. This can be attributed to the fact that they are accreting material more efficiently. Lastly, we compare the AT20G survey with the S-cubed semi-empirical (S3-SEX) models and conclude that the S3-SEX models need refining to correctly model the compact cores of AGN. The AT20G survey provides the ideal sample to do this. ",
        "introduction": "\\label{intro}  The formation and evolution of AGN, and the role they play in galaxy evolution, is currently one of the most active fields in astronomy. The advent of large wide-area surveys across a multitude of frequencies, from radio to gamma--rays, has provided large, homogenous samples allowing us to study the different AGN populations in a statistically significant manner.  The radio sky at frequencies of 5\\,GHz and below has been well explored, largely due to surveys such as the 843\\,MHz Sydney University Molonglo Sky Survey (SUMSS; \\citealt{sumss}), 1.4\\,GHz NRAO VLA Sky Survey (NVSS; \\citealt{nvss}) and the 5\\,GHz Parkes-MIT-NRAO (PMN; \\citealt{pmn}) and Green Bank 6\\,cm (GB6; \\citealt{gb6}) surveys. The Australia Telescope 20\\,GHz (AT20G) Survey provides an unprecedented view of the high-frequency radio sky. This blind survey was carried out from 2004--2008 on the Australia Telescope Compact Array (ATCA) and covers the entire southern sky excluding $|b|<1.5^{\\circ}$. The AT20G catalogue consists of 5890 sources above 40\\,mJy \\citep{at20g}. To obtain reliable spectral index information, observations at 4.8 and 8.6\\,GHz were carried out within 1 month of the 20\\,GHz observations for most sources south of $\\delta=-15^{\\circ}$.  Selecting sources at 20\\,GHz provides a unique sample that is dominated by active galaxies which derive most of their radio emission from supermassive black holes (SMBH)\\footnote{Besides a small number of galactic sources (namely PN and HII regions) the only known non-AGN in the AT20G survey is NGC253; a nearby galaxy which hosts a strong nuclear starburst \\citep{ngc253}.}. On the other hand, low-frequency selected samples contain a mixture of AGN and star-forming galaxies \\citep{mauch07}. Furthermore, for AGN that are observed at lower frequencies, the dominant source of emission is the large-scale radio lobes which have been built up over much longer timescales, whereas observing at high radio frequencies provides insight into the most recent activity in the central SMBH.  Understanding the high-frequency radio population is vital to understanding the growth and evolution of AGNs and provides complementary data to their low-frequency counterparts. An analysis of the radio properties of AT20G show that it is a very different radio population from the low-frequency selected radio population. \\citet{at20ganalysis} find that the catalogue is dominated by flat-spectrum sources, particularly at bright fluxes where 81\\% of the sample has flat radio spectra (defined as $\\alpha^{20}_8 > -0.5$ where $S_{\\nu}\\propto\\nu^{\\alpha}$) which drops of to 60\\% for fluxes below 100\\,mJy. A recent review by \\citet{dezotti10} provides a comprehensive view of the radio properties of AGN, including both low and high-frequency information.  Whilst the radio data provides a unique sample of objects, these data alone are insufficient to constrain models of radio source properties and the evolution of radio galaxies. Additional optical data, particularly spectroscopic information, is vital in understanding the physical properties of the central black hole.  There have been a number of other surveys carried out at high radio frequencies. These include; the Wilkinson Microwave Anisotropy Probe (WMAP), which covers the whole sky at 22\\,GHz down to a 1\\,Jy flux limit, the Planck Early Release Compact Source Catalogue \\citep{planckcat, planckanalysis}, an all-sky survey at 30\\,GHz, with a flux limit of 0.5\\,Jy and the 9th and 10th Cambridge surveys (9C; \\citealt{9c}, and 10C; \\citealt{10c}) carried out at 15\\,GHz that cover a much smaller area, but observe down to fainter flux limits (5.5\\,mJy and 1.0\\,mJy respectively). However, with the exception of \\citet{bolton04} who followed up sources in the 9C sample at both optical and radio frequencies, there has not been much focus on the optical properties of these surveys.  It has been known for sometime that radio galaxies can show a range of optical spectral signatures. Some galaxies display strong emission lines whilst others have very weak or no emission features (\\citealt{hine+longair} and more recently \\citealt{sadler02}). In recent years these differences have been attributed to different accretion modes \\citep{best05, hardcastle07}. `Cold-mode' accretion occurs when cold gas is accreted onto the central SMBH, which gives rise to the traditional AGN picture of an accretion disk surrounded by a dusty torus, a broad-line region close to the accretion disk and narrow-line region further out \\citep{antonucci93}. Such objects show high-excitation lines in the optical spectrum, either narrow or broad depending on the alignment of the AGN to our line of sight. `Hot-mode' accretion occurs when hot-gas is being accreted, resulting in a radiatively inefficient accretion disk that observationally shows none of the conventional signs of AGN activity in the optical regime. It is theorised that cold-mode (also referred to as `quasar-mode') accretion occurs as a result of mergers which supply a reservoir of cold gas, whilst hot-mode (also known as `radio-mode') accretion is the result of hot gas from the surrounding ISM falling into the central core \\citep{croton06}.  It was thought that these different accretion modes could account for the observed dichotomy of radio-loud AGN \\citep{fanaroff}. However, it has been shown that in general the radio properties of AGN observed at 1.4\\,GHz are completely independent of the optical emission line properties \\citep{best05}, but are well correlated with the stellar mass of the galaxy (see also \\citealt{mauch07}). By studying the optical properties of a high-frequency radio sample we aim to investigate the link between the 20\\,GHz emission (predominantly from the core of the AGN) and the optical properties since both are generated on similar size and time scales.  Additional optical and redshift information of high-frequency radio sources also provides valuable information for both the Fermi and ESA's Planck missions, which rely on multiwavelength data at higher angular resolution to correctly identify their sources. It has already been shown that there is large overlap between the AT20G survey and both the Fermi Large Area Telescope First Source Catalog ({\\it Fermi} 1FGL; \\citealt{fermi1FGL}) and the Early Release Compact Source Catalog \\citep{planckcat}; see \\citet{fermiat20g} and \\citet{paco} respectively. The postional uncertainty of these instruments require multiwavelength information to accurately identify the correct counterpart. The positional accuracy of AT20G ($\\sim 1$\\,arcsec) allows us to identify the correct optical source and redshift, providing valuable information for these surveys.  The primary goal of this paper is to present accompanying optical and redshift data for the AT20G survey, but we also present some of the basic population statistics of a high radio-frequency selected survey. Section~\\ref{makingids} details how the optical identifications were made and Section~\\ref{redshifts} discusses how redshifts were found. Section~\\ref{specclass} describes the spectral properties as defined from the optical spectrum and Section~\\ref{catalogue} describes the resulting AT20G--optical catalogue. In Section~\\ref{results} we present results from the radio--optical analysis. This is followed by a comparison of the AT20G catalogue with the S-cubed semi-empirical models \\citep{s3-sex}, which make predictions of the 18\\,GHz sky, in Section~\\ref{skads}. We conclude in Section~\\ref{conclusions}. Throughout this paper we use the following cosmological parameters: $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$ and $\\Omega_{\\Lambda}=0.73$ \\citep{wmap7}. ",
        "conclusions": "\\label{conclusions}  The primary goal of this paper is to present a catalogue of the optical identifications and redshifts for the Australia Telescope 20\\,GHz (AT20G) survey. The positional accuracy ($\\sim 1$\\,arcsec) of AT20G sources, and the dominance of nuclear point sources meant that the majority of sources were identified using an automated process. Optical identifications were made using the SuperCOSMOS Science Archive with 78.5\\% of AT20G sources (with $|b|>10^{\\circ}$) having optical counterparts, much larger than is seen at lower radio frequencies (typically 25--30\\%, \\citealt{bock99}). The optical identification rate increases as a function of 20\\,GHz flux density, predominately due to the increase in `stellar' optical IDs with increasing flux density. On the other hand, sources with `galaxy' counterparts (as determined by SuperCOSMOS based on the optical morphology) decreased with increasing flux density. Whilst the `stellar' sources are the dominant population in the AT20G survey, we predict that if we were to probe deeper at 20\\,GHz the galaxy population would become dominant.  Redshifts were found from either the 6dF Galaxy Survey \\citep{6df} or from the literature using NED. We also present 144 new redshifts that were obtained from poor-weather backup programs from a range of facilities. A total of 30.9\\% of the AT20G survey have redshift information. The redshift completeness is quite high for the bright AT20G sources (91.5\\% complete for sources brighter than 1\\,Jy and 85.3\\% for sources above 500\\,mJy), but this drops off rapidly towards fainter sources.  Objects that have an optical spectrum available (i.e. from either the 6dF survey or additional observations) formed the `AT20G spectroscopic sample' and this sample is used to study the emission line properties of the sample. This is achieved by dividing the sample into `hot' or `cold' mode accretors, where sources with strong emission lines are classified as undergoing cold-mode accretion and sources with weak or no emission lines present were classified as hot-mode. This naming convention arises from the fact that cold-mode sources primarily accrete cold gas, forming a radiatively efficient accretion disk which fits the conventional AGN unification models. On the other hand, hot-mode sources accrete hot gas which leads to an inefficient accretion disk and none of the observational criteria normally associated with optical or X-ray selected AGN.  It is found that cold-mode accreting sources have a steeper median radio spectral index, suggesting that the radio jets of these sources dominate the emission (rather than the cores), consistent with the idea that these sources are more efficient accretors. It is also found that the cold-mode accretors show different orientation effects; narrow emission line objects (Ae sources) have a larger fraction of steep-spectrum radio sources, while objects with broad emission lines (AeB sources) are dominated by flat-spectrum sources as predicted by unification models \\citep{antonucci93, urry+padovani}. On the other hand, the hot-mode accretors show no distinct orientation effects.  Due to the high-frequency selection, the AT20G sample is dominated by flat-spectrum sources, as are many of the subsamples. AT20G sources with optical IDs, the redshift sample and the spectroscopic sample are all dominated by flat-spectrum sources, yet for the subset of sources that do not have an observed optical counterpart (blank fields) the split into flat and steep spectrum sources is approximately 50/50. This implies that many of these blank fields exhibit narrow emission lines (falling into the Ae spectral class) as this is the only population of sources dominated by steep-spectrum sources. Some of these blank fields also form a small sample (36 objects) of high-frequency ultra-steep spectrum (USS) sources, listed in Appendix \\ref{appendixb}. As the peak of the radio SED is thought to move down in frequency as the source ages, it is likely that these are the progenitors of the USS sources selected at lower frequencies.  Studying the radio luminosity vs. absolute magnitude distribution reveals that, for a given radio luminosity, the cold-mode accretors are more likely to reside in smaller host galaxies whilst the hot-mode accreting sources generally live in more massive galaxies. Objects accreting in the cold-mode also have higher radio luminosities for the same absolute magnitude than their hot-mode counterparts, providing further evidence that the cold-mode accreting sources produce more luminous radio jets and lobes.  Lastly, we compared the AT20G survey with the S-cubed semi-empirical simulations (S3-SEX). This was achieved by selecting sources from the S3-SEX database to be brighter than 100\\,mJy (the AT20G completeness limit) at 18\\,GHz. The flux distributions of the S3-SEX and AT20G samples are in agreement with each other, but the spectral index distributions are remarkably different, implying that the models used to determine the S3-SEX core fluxes need refining. This comparison highlights the need for large samples of high-frequency radio sources provided by the AT20G survey.  We continue to compile spectroscopic information on AT20G sources with the aim of forming a complete sample to further study the properties of high-frequency radio sources and how this population evolves over cosmic time."
    },
    "1107/1107.0721_arXiv.txt": {
        "abstract": "In a recent paper~\\cite{DynamicalDM1}, we introduced ``dynamical dark matter,'' a new framework for dark-matter physics, and outlined its underlying theoretical principles and phenomenological possibilities.  Unlike most traditional approaches to the dark-matter problem which hypothesize the existence of one or more stable dark-matter particles, our dynamical dark-matter framework is characterized by the fact that the requirement of stability is replaced by a delicate balancing between cosmological abundances and lifetimes across a vast ensemble of individual dark-matter components. This setup therefore collectively produces a time-varying cosmological dark-matter abundance, and the different dark-matter components can interact and decay throughout the current epoch.  While the goal of our previous paper was to introduce the broad theoretical aspects of this framework, the purpose of the current paper is to provide an explicit model of dynamical dark matter and demonstrate that this model satisfies all collider, astrophysical, and cosmological constraints.  The results of this paper therefore constitute an ``existence proof'' of the phenomenological viability of our overall dynamical dark-matter framework, and demonstrate that dynamical dark matter is indeed a viable alternative to the traditional paradigm of dark-matter physics.  Dynamical dark matter must therefore be considered alongside other approaches to the dark-matter problem, particularly in scenarios involving large extra dimensions or string theory in which there exist large numbers of particles which are neutral under Standard-Model symmetries. ",
        "introduction": "The nature of what constitutes the non-baryonic dark matter in our universe remains one of the most fundamental mysteries in particle physics~\\cite{DMReviews}. The most precise measurements of the relic abundance of this dark matter to date are those derived from WMAP data~\\cite{WMAP}, which yield a value \\begin{equation} \\Omega_{\\mathrm{CDM}} h^2 ~=~ 0.1131 \\pm 0.0034~, \\label{eq:OmegaWMAP} \\end{equation} where $h \\approx 0.72$ is the Hubble constant.  Beyond this, we know very little about the properties of this dominant constituent of the matter density in our universe, save that its interactions with the fields of the Standard Model (SM) are extremely weak. One of the reasons why the nature of the dark matter remains so elusive is its apparent stability. Observational constraints on the lifetime $\\tau_\\chi$ of any decaying dark-matter candidate $\\chi$ are quite stringent.  Indeed, for any particle with a relic abundance $\\Omega_\\chi \\sim \\OmegaDM$, current limits~\\cite{DMDecayChenKamionkowski} from cosmic microwave background (CMB) measurements, \\etc, require that \\begin{equation} \\tau_\\chi ~\\gtrsim~  10^{26}\\mathrm{~s}~. \\label{eq:DecayingDMLifetimeLimit} \\end{equation} For this reason, most models of the dark sector posit the existence of a single dark-matter particle (or, in the case of certain multi-component dark-matter scenarios~\\cite{MultiCompDMBlock,Winslow}, a small number of such particles) which is either absolutely stable (with that stability usually conferred by some additional symmetry, such as R-parity in supersymmetric models, KK-parity~\\cite{KKParity} in universal extra dimensions~\\cite{Antoniadis,DDGLargeED,UED}, or T-parity~\\cite{TParity} in little-Higgs theories~\\cite{LittleHiggs}), or else sufficiently long-lived as to satisfy the bound in Eq.~(\\ref{eq:DecayingDMLifetimeLimit}).  Indeed, the phenomenological consequences of dark-matter decays in models with unstable dark-matter candidates~\\cite{DecayingDM} can be quite significant.  Recently, an alternative framework for addressing the dark-matter question has been proposed~\\cite{DynamicalDM1}.  In this so-called ``dynamical dark matter'' paradigm, the dark sector comprises not one or merely a few particle species, but rather a vast ensemble of different fields $\\phi_i$, each of which contributes only a fraction $\\Omega_i$ of the total dark-matter relic abundance $\\OmegaDM$.  None of these fields is presumed to be absolutely stable, and thus a non-zero decay width $\\Gamma_i$ is associated with each field. However, in this framework, the individual relic abundances of the $\\phi_i$ fields are presumed to be generated in such a way that the most stable members of that ensemble are the most abundant.  By contrast, the abundances of the more unstable members are suppressed according to the size of their decay widths.  It is this balancing between $\\Gamma_i$ and $\\Omega_i$ which makes it possible for the phenomenological constraints relating to the effects of dark-matter decays to be satisfied.  In Ref.~\\cite{DynamicalDM1}, we focused on the model-independent aspects of our dynamical dark-matter framework, discussing its broad theoretical properties, without any detailed phenomenological analysis or comparison with data.  By contrast, in this work, as a ``proof of concept,'' we provide an explicit model of dynamical dark matter. In this model, the fields which collectively constitute the dynamical dark-matter ensemble are the KK excitations of a light axion-like field propagating in the bulk of a spacetime with one or more large, flat extra dimensions.  In this model, the fields of the SM, as well as the gauge fields associated with some additional, non-Abelian gauge group $G$ which confines at a scale $\\Lambda_G$, are taken to be localized on a four-dimensional subspace of that bulk.  The axion field is assumed to couple to the gauge fields of $G$ (and also potentially to one or more of the SM fields) via non-renormalizable operators suppressed by some effective, four-dimensional cutoff scale $\\fhatX$. We shall demonstrate that within this setup, the resulting ensemble of axion KK modes naturally satisfies all applicable observational constraints on dark-matter decays --- even if the stability of this ``dark tower'' is entirely unprotected.  Another advantage of this particular dynamical dark-matter model is that it is not only phenomenologically viable, but also theoretically well motivated. In any theory in which the SM fields reside on a brane, the KK excitations of any bulk field are, from the point of view of the four-dimensional theory, massive particles neutral under the SM gauge group.  Thus, were it not for the lack of a stabilizing symmetry, any of these particles would be a natural candidate for dark matter.  However, our results demonstrate that a lack of stability is not an insurmountable impediment to such fields serving as dark matter {\\it collectively}, rather than individually. This is an exciting prospect, for it provides a novel way of addressing the dark-matter question in theories with extra dimensions. Furthermore, our model also demonstrates that realizing a viable dynamical dark-matter ensemble does not require an overly complicated dark sector, a large number of independent mass scales, or an excessive degree of fine-tuning. Indeed, the model presented here involves only three independent physical scales: the effective four-dimensional cutoff scale $\\fhatX$, the confinement scale $\\Lambda_G$, and the compactification scale $M_c$.  Together, these three scales determine the mass spectrum and decay properties of the entire ensemble.  The outline of this paper is as follows. In Sect.~\\ref{sec:AxionsInED}, we briefly review the formalism for discussing axions and axion-like fields, beginning with the standard, four-dimensional case and then moving on to the generalized, five-dimensional bulk-axion case.  In Sect.~\\ref{sec:AxionDecayRates}, we calculate the decay widths of the KK modes of such a bulk axion-like field and investigate how these decay widths scale with the mass of the mode.  In the process, we show that the decays of the lighter modes to SM fields experience a natural suppression, but that such decays nevertheless dominate over decays to other, lighter bulk fields in the theory. In Sect.~\\ref{sec:Production}, we examine the various mechanisms through which a population of axion modes may be generated in the early universe, and demonstrate that the abundances generated for those modes by misalignment production are indeed balanced against their decay widths in precisely the manner required for dynamical dark matter.  In Sect.~\\ref{sec:Abundances}, we examine the collective properties of the ensemble of axion KK modes.  We show that such an ensemble can collectively reproduce the observed dark-matter relic density given in Eq.~(\\ref{eq:OmegaWMAP}), and that it possesses the appropriate equation of state to be regarded as dark matter.  In Sect.~\\ref{sec:Constraints}, we summarize the experimental, astrophysical, and cosmological constraints on scenarios involving bulk axions in large, flat extra dimensions.  We demonstrate that these constraints can be satisfied in a model which also simultaneously yields the correct total relic abundance --- in other words, that our model truly constitutes a viable model of dynamical dark matter.  Finally, in the Conclusions, we summarize the results of the previous sections and discuss several further directions for future investigation.  As we have indicated, this paper is the second part of a two-part series that began with Ref.~\\cite{DynamicalDM1}. Consequently, we shall assume that the reader is familiar with the ideas, notation, and conventions established in Ref.~\\cite{DynamicalDM1} in what follows. ",
        "conclusions": "}   The aim of this paper has been to present an explicit realization of the dynamical dark-matter framework presented in Ref.~\\cite{DynamicalDM1}. To that end, we have shown that an ensemble consisting of the KK excitations of a light, axion-like field can indeed provide such a realization.  Indeed, we have shown that despite the fact that the masses, decay widths, and relic abundances of all of these particles are controlled by only three dimensionful parameters, the ensemble to which they give rise is simultaneously able to reproduce the observed value of $\\OmegaDM$ and satisfy all applicable constraints from laboratory experiments, astrophysics, and cosmology. As such, this model provides a ``proof of concept'' for dynamical dark matter as a viable alternative framework for dark-matter physics.  In addition, it also provides a method of addressing the dark-matter question which does not require the introduction of any additional stabilizing symmetry.  Many qualifications, extensions, and possible generalizations of our dynamical dark-matter framework were discussed at the end of Ref.~\\cite{DynamicalDM1}; here, we shall restrict our attention to five points which are specific to the bulk-axion model presented in this paper.  \\begin{itemize}  \\item First, in this work, we have made use of the rapid-turn-on approximation in Eq.~(\\ref{eq:Heaviside}) in calculating the relic abundances of the $a_\\lambda$.  As discussed in Sect.~\\ref{sec:Production}, this approximation is well motivated, since the instanton-generated mass term $\\mX(T)$ falls rapidly with temperature when $T\\gtrsim \\Lambda_G$.  Furthermore, the primary results of this paper are essentially insensitive to this approximation.  This is because the fields which contribute significantly to $\\OmegaDM$ in regions of parameter space which yield a realistic dark-matter relic abundance begin oscillating only well after $\\mX(T)$ has already settled into its constant, late-time value.  However, the relic abundance of any field which begins oscillating before $\\mX(T)$ takes this late-time value will, in general, depend on the details of how this mass evolves in time.  The quantitative effect on the abundance of a single field has long been appreciated~\\cite{Turner}, but in our model, the effects are more complicated and more subtle because we have a coupled system of mixed scalars with different masses and therefore different oscillation times.  It would be interesting to examine how a more rigorous treatment of the turn-on of $\\mX(T)$ would affect $\\Omegatot$ and $\\eta$ in situations in which these quantities are sensitive to the time-dependence of this brane-mass term.  Such a study would have important implications for more general scenarios involving other kinds of light bulk scalars.  Indeed, the relationship between the size of the brane-mass term for such scalars and the time at which that brane mass is dynamically generated may differ significantly from the relationship which holds for axions.  \\item Second, as alluded to in Sect.~\\ref{sec:Constraints}, it may be possible to further test or constrain the parameter space of bulk-axion models of dynamical dark matter in a number of ways.  We have already mentioned one potential constraint which derives from limits on high-energy photons resulting from the decays of axions produced in supernovae~\\cite{GiannottiDuffy}. Other considerations also merit investigation.  For example, a detailed analysis of the limits imposed by BBN on scenarios involving multiple decaying fields with different lifetimes and abundances could provide important constraints on dynamical dark-matter models in general.  In addition, other considerations, such as limits on mass loss and decreases in the dark-matter density in the halos of dwarf galaxies~\\cite{DwarfGalaxyDecayingDM}, could also be used to constrain dynamical dark-matter models.  Indeed, while a number of standard constraints on individual unstable relic particles in the early universe have been revisited in a dynamical dark-matter context~\\cite{DynamicalDM3}, it would be interesting to see how other constraints would apply in this context as well.   \\item Third, we note that we have not specified a particular model of inflation as part of the cosmological context for our model.  Indeed, other than requiring a low reheating temperature $\\TRH \\sim \\mathcal{O}(\\mathrm{MeV})$, we have remained largely agnostic about the details of the inflationary model, the form of the inflaton potential, or even the scale $H_I$.  For the most part, our model does not depend on these particulars.  However, certain consistency conditions do place meaningful restrictions on the set of inflationary scenarios with which our model is compatible.  One such condition can be derived from the fact that vacuum fluctuations during inflation generically give rise to a background value $\\langle \\phi^2 \\rangle \\approx H_I^3 t_I/4\\pi^2$ for any scalar $\\phi$ with a mass $m_\\phi \\ll H_I$, where $t_I$ is the duration of inflation.  This implies that the relationship between the mass $\\lambda$ and initial energy density $\\rho_\\lambda$ in Eq.~(\\ref{eq:RhoInitCondits}) in our model is truly valid only for the lighter $a_\\lambda$ in a given tower --- \\ie, those for which $\\theta^2 A_\\lambda^2 \\fhatX^2 \\gtrsim H_I^3 t_I/4\\pi^2$.  By contrast, any heavier $a_\\lambda$ which still satisfy $\\lambda \\ll H_I$ receive the leading contributions to their background values from vacuum fluctuations during inflation, and thus effectively acquire an initial abundance $\\rho_\\lambda \\sim \\lambda^2 H_I^3 t_I$.  In typical scenarios, we expect $H_I t_I \\approx N_e \\sim \\mathcal{O}(60)$, where $N_e$ is the number of $e$-foldings of inflation.  The results for $\\Omegatot$ derived in Sect.~\\ref{sec:Abundances} therefore remain consistent, provided that $\\fhatX^2 \\gg H_I^2$.  Indeed, since $\\fhatX \\sim 10^{14} - 10^{15}$~GeV within the preferred region of parameter space specified in Eq.~(\\ref{eq:PreferredRegion}), we see that $\\fhatX \\gg H_I$ is certainly not inconsistent with our model and is in fact even expected.  However, this condition on $H_I$ has non-trivial implications for inflationary models. While a low scale for $H_I$ is certainly not excluded (see, \\eg, Refs.~\\cite{LTRAxionsKamionkowski,lowscaleinflation}), extremely small values of $H_I$ tend to be rather non-generic~\\cite{TensorToScalarNonGeneric} among typical classes of inflationary potentials, and thus require either substantial tuning or careful construction.  Indeed, any consistent inflationary model of this sort must give rise to density fluctuations on a scale consistent with constraints from CMB data~\\cite{WMAP}, such as those on the spectral index $n_s$, and must also satisfy other observational constraints.  The development of explicit inflationary scenarios of this sort is therefore an interesting topic for future investigation.  \\item Fourth, we note that while we have chosen in this paper to focus on the case in which the ensemble of fields reproducing $\\OmegaDM$ are the KK excitations of a bulk axion field, such a field is by no means unique in possessing the characteristics necessary to give rise to such an ensemble.  Indeed, as discussed in Ref.~\\cite{DynamicalDM1}, much of the analysis presented here pertains to any light bulk scalar for which a mass term is dynamically generated via its interactions with brane-localized fields.  Furthermore, for a generic bulk scalar, the relationship between the time at which this mass term is dynamically generated and the magnitude of this mass term itself may differ from that which relates $t_G$ and $\\mX$ for a bulk axion.  As a result, much more freedom may exist for constructing viable models within the dynamical dark-matter framework.  For example, light moduli could also, in principle, provide a viable model of dynamical dark matter.  \\item Finally, we emphasize that the presence of additional axion-like fields is fairly generic, and perhaps even expected, in many theoretically motivated scenarios for physics beyond the Standard Model (see, \\eg, Ref.~\\cite{WittenStringAxion}).  Moreover, it has even been argued that many of these axion-like fields are likely to be light~\\cite{StringAxiverse}. Thus, the discovery of a vast ensemble of axion-like particles could provide important insight into what physics looks like at high scales. Indeed, if many of these axions have relatively small masses, we find ourselves in the intriguing situation in which most of the matter in the universe is simultaneously both light and dark.  \\end{itemize}  Our goal in this work has been to provide an existence proof for dynamical dark matter --- \\ie, to provide a model in which lifetimes are balanced against abundances in such a way that the ensemble of dark-matter particles successfully reproduces $\\OmegaDM$ while at the same time satisfying all phenomenological constraints.  As we have seen in this paper, our bulk-axion model indeed passes this test.  In one sense, our model does so in the most interesting way possible: with $y\\ll 1$ (signifying that our tower of axion KK modes is highly mixed) and with a tower fraction $\\eta$ which is significantly different from zero.  In another sense, however, this model is fairly conservative: those modes which contribute most to $\\Omegatotnow$ turn out to be rather long-lived, and likewise our numerical result for $w_\\ast$ within the preferred region of parameter space turns out to be rather close to zero.  Indeed, at first glance, one might suspect that these latter properties are in fact generic for dynamical dark-matter models, or even that such models are therefore really no different from traditional dark-matter models in terms of their abundance and stability requirements.  This is not the case, however, for the balancing of lifetimes against abundances --- which is the hallmark of the dynamical dark-matter framework --- is precisely why this framework does {\\it not} require such a degree of stability, much less the existence of a stabilizing symmetry. While certain accidental features of our bulk-axion model result in a preferred region of parameter space which is somewhat conservative, we emphasize that these features are not generic even to theories with bulk scalars, much less realistic dynamical dark-matter models as a whole.  Note, for example, that a particular relationship exists in bulk-axion models between the mass $\\lambda$ of a given KK mass eigenstate $a_\\lambda$, the strength of its effective coupling to SM fields, and the overall magnitude of its relic abundance $\\Omega_\\lambda$ through the dependence of these quantities on $\\fhatX$.  Even for other bulk scalars (\\eg, moduli), these relationships do not necessarily hold.  There is therefore no reason to expect dynamical dark-matter models based around such fields to be as conservative as the axion model we have presented here.  In this connection, there is an even more important point that deserves emphasis.  In dynamical dark-matter scenarios, we have no single characteristic decay width $\\Gamma$ nor abundance $\\Omega$, but rather an entire {\\it spectrum}\\/ of widths $\\Gamma_\\lambda$ and abundances $\\Omega_\\lambda$.  This therefore begs the fundamental question:  if our ``proof of concept'' model presented here is to be viewed as somewhat conservative, how far from the conservative limit can we go?  At first glance, one might try to answer this question by attempting to determine, for each time $t$ during the evolution of the universe, the maximum abundance $\\Omega_{\\mathrm{max}}(t)$ that a given component in a dark-matter ensemble may have if it has a lifetime $\\tau \\sim t$.  In other words, given the entirety of the cosmological constraints from BBN, CMB distortions, \\etc, there exists a {\\it function}\\/ $\\Omega_{\\mathrm{max}}(\\Gamma)$ which describes the maximum abundance any dark-matter constituent may have as a function of its decay width.  It might therefore seem that knowledge of this function would uniquely determine the full range of possibilities inherent in our dynamical dark-matter framework.  Such an approach to answering our fundamental question is, in a sense, already a departure from the usual manner of approaching dark-matter physics. However, even the notion of such a function $\\Omega_{\\rm max}(\\Gamma)$ relies too strongly on a single-particle perspective. One of the critical features of our dynamical dark-matter framework is that it involves a vast {\\it ensemble}\\/ of dark-matter components. Some of these components might decay earlier in cosmological evolution, while others might decay later. As a result, the maximum abundance that a given component may have if it decays on a characteristic time scale $\\tau$ will itself be directly affected not only by the abundances of all of the other components with earlier characteristic decay times $\\tau' < \\tau$, but even the components with $\\tau' > \\tau$.  Moreover, as we have seen, most phenomenological constraints on dark-matter decays are sensitive not merely to what happens at a specific moment in time, but to the integrated effects of such decays over a broad range of time scales. In other words, our dynamical dark-matter framework teaches us that astrophysical and cosmological constraints do not lead to a single function $\\Omega_{\\rm max}(\\Gamma)$, but rather a more subtle set of intertwined constraints on lifetimes and abundances across our entire dark-matter ensemble as a whole.  Clearly, this issue has not been studied in any detail in the literature. However, it is readily apparent that this is indeed the only proper way in which one should express constraints on particle decays from a generic dark sector.  Viewed from this perspective, then, the existence of even one viable dynamical dark-matter model --- no matter how ``conservative'' it might be --- gives us strong motivation to re-examine cosmological and astrophysical constraints within this framework. Indeed, it is only in this way that we will be able to fully explore our dynamical dark-matter framework, and understand its full range of phenomenological possibilities."
    },
    "1107/1107.0517_arXiv.txt": {
        "abstract": "We examine the formation of groups of multiple supermassive black holes (SMBHs) in gas-poor galactic nuclei due to the high merger rate of galaxies at high redshifts.  We calculate the relative likelihood of binary, triple, and quadruple SMBH systems, by considering the timescales for relevant processes and combining merger trees with N-body simulations for the dynamics of stars and SMBHs in galactic nuclei.  Typical haloes today with mass $M_0\\approx 10^{14}$ M$_\\odot$ have an average mass $M_{z=6}=5\\times 10^{11}$ M$_\\odot$ at $z\\sim 6$, while rare haloes with current mass $M_0\\gtrsim 10^{15}$ M$_\\odot$ have an average mass $M_{z=6}=5\\times 10^{12}$ M$_\\odot$ at that redshift.  These cluster-size haloes are expected to host single galaxies at $z\\sim 6$.  We expect about 30\\% galaxies within haloes with present-day mass $M_0\\approx 10^{14}$ M$_\\odot$ to contain more than two SMBHs at redshifts $2\\lesssim z\\lesssim 6$.  For larger present-day haloes, with $M_0\\gtrsim 10^{15}$ M$_\\odot$, this fraction is almost 60\\%.  The existence of multiple SMBHs at high redshifts can potentially explain the mass deficiencies observed in the cores of massive elliptical galaxies, which are up to $5$ times the mass of their central BHs.  Multiple SMBHs would also lead to an enhanced rate of tidal disruption of stars, modified gravitational wave signals compared to isolated BH binaries, and slingshot ejection of SMBHs from galaxies at high speeds in excess of $2000$ km s$^{-1}$. ",
        "introduction": "\\label{sec:intro}  Most local galaxies host supermassive black holes (SMBHs) at their centres \\citep{1998Natur.395A..14R, 2005SSRv..116..523F}.  The SMBH mass $M_{\\rm bh}$ is correlated with properties of the spheroidal nucleus of the host galaxy, such as velocity dispersion \\citep{2000ApJ...539L...9F, 2000ApJ...539L..13G, 2002ApJ...578...90F, 2009ApJ...698..198G} and luminosity \\citep{1998AJ....115.2285M, 2002MNRAS.331..795M, 2003ApJ...589L..21M, 2009ApJ...698..198G}. Detection of bright quasars at redshifts $z\\gtrsim 6$ \\citep{2001AJ....122.2833F, 2011Natur.474..616M} suggests that SMBHs with masses as high as $\\sim 2\\times 10^9$ M$_\\odot$ already existed at $z\\sim 7$.  In the standard $\\Lambda$CDM cosmological model, growth of galaxies is hierarchical and galaxy mergers are expected to be particularly frequent at redshifts $z\\sim 6$--$20$.  As galaxies merge, their central SMBHs can grow through coalescence and accretion of gas.  It is commonly postulated that SMBHs at lower redshifts grew out of seed black holes (BHs) in the first galaxies \\citep{1994ApJ...432...52L, 1995ApJ...443...11E, 2000MNRAS.311..576K, 2001ApJ...558..535M, 2003ApJ...596...34B, 2003ApJ...582..559V, 2006ApJS..163....1H, 2009ApJ...696.1798T}.  Existing merger tree models are based on the assumption that any binary black hole system, which inevitably forms in a galaxy's merger history, coalesces on a short time-scale.  However, the evolution of SMBH binaries is a complex open problem and it is unclear if a binary can merge within a Hubble time \\citep {2005LRR.....8....8M}.  One expects that during a merger event of two galaxies, the dynamics of their constituents SMBHs would proceed in three stages \\citep{1980Natur.287..307B}.  In the first stage, the SMBHs sink to the centre of the gravitational potential of the merger remnant by dynamical friction and form a gravitationally bound binary.  The newly-formed binary continues to lose energy and angular momentum through its global gravitational interaction with many stars until the separation between the SMBHs reduces to a value at which the dominant mechanism of energy loss is the 3-body interaction between the binary and individual stars.  This is the second stage of the binary's evolution, and is known as the `hard stage.'  The precise definition of a hard SMBH binary varies in the literature, but it is commonly assumed that the binary becomes hard when its semi-major axis $a$ reaches a value given by \\citep{2002MNRAS.331..935Y} \\begin{equation} a\\approx a_h\\equiv\\frac{Gm}{4\\sigma^2}=2.8\\left(\\frac{m}{10^8 \\mathrm{M}_\\odot}\\right)\\left(\\frac{200\\;\\mathrm{km}\\,\\mathrm{s}^{-1}}{\\sigma}\\right)^2 \\mathrm{pc}, \\label{eqn:ahard} \\end{equation} where stars in the galactic nucleus are assumed to have a one-dimensional velocity dispersion $\\sigma$, and $m$ denotes the mass of the lighter SMBH.  Finally, once the SMBH separation decreases to a small-enough value, gravitational wave emission becomes the dominant mode of energy loss and the SMBHs coalesce rapidly.  This is the third stage of the SMBH binary evolution.  The value of semi-major axis $a$ at which the coalescence time scale due to gravitational wave emission alone is $t$ is given by \\citep{peters1964, 2010PhRvD..81d7503L} \\begin{equation} a(t)\\equiv a_\\mathrm{gw}(t)=4.3\\times 10^{-3}\\left(\\frac{t}{10^{5}\\mathrm{yr}}\\right)^{1/4}\\left(\\frac{M}{2\\times 10^{8}\\mathrm{M}_\\odot}\\right)^{3/4}\\!\\mathrm{pc}, \\label{eqn:agw} \\end{equation} where $M$ is the total mass of the binary, and we have considered two SMBHs with mass $10^8$ M$_\\odot$ each on a circular orbit (with shorter time scale at increasing eccentricity).  Gravitational wave emission takes over as the dominant mode of energy loss when $a=a_\\mathrm{gw}(t_h)$, where $t_h$ is the hardening time scale.  Among these three stages of evolution of an SMBH binary, the largest uncertainty in the binary's lifetime originates from the hard stage, which can be the slowest of the three stages since the binary quickly ejects all low angular momentum stars in its vicinity, thus cutting off its supply of stars.  This is known as the ``final parsec problem'' \\citep{2003AIPC..686..201M}.  For example, \\citet{2002MNRAS.331..935Y} studied coalescence of SMBH binaries in a sample of galaxies observed by \\citet{1997AJ....114.1771F} and found that spherical, axisymmetric or weakly triaxial galaxies can all have long-lived binary SMBHs that fail to coalesce.  Similarly, \\citet{2005LRR.....8....8M} found that the time spent by a binary is less than $10^{10}$ yr only for binaries with very low mass ratios ($\\lesssim 10^{-3}$).\\footnote{However, for such low mass ratios the time taken by the lighter black hole to reach the galactic nucleus due to dynamical friction is itself expected to exceed the Hubble time.} Furthermore, \\citet{2005LRR.....8....8M} showed that a binary may not be able to interact with all the stars in its loss cone, thereby increasing the time spent in the hard stage even further; they found that in a nucleus with a singular isothermal sphere stellar density profile, an equal-mass binary will stall at a separation of $a\\approx a_h/2.5$, where we have defined $a_h$ in Equation (\\ref{eqn:ahard}). The final separation is expected to be even higher for galaxies with shallower density profiles.  Several ways have been discussed in the literature to efficiently extract energy and angular momentum from a hard SMBH binary and overcome the final parsec problem.  An example is the work by \\citet{2002ApJ...567L...9A}, who suggested that gas can catalyse the coalescence of a hard SMBH binary by serving as an effective sink for the binary's angular momentum.  In particular, they found that a binary with a separation of $0.1$ pc embedded in a gaseous accretion disk would merge in $10^7$ years without significant enhancement in the gas accretion rate.  Similarly, \\citet{2004ApJ...607..765E, 2005ApJ...630..152E} found that in SPH simulations, clouds of hot gas ($T_\\mathrm{gas}\\approx T_\\mathrm{virial}$) can induce decay of orbits of embedded binary point masses due to gravitational drag.  A caveat to these studies is that feedback from gas accretion onto the SMBHs can remove the rest of the gas from the merger remnant before the binary coalesces.  However, stellar dynamical processes could also accelerate binary coalescence, without gas.  For example, \\citet{2004ApJ...606..788M} considered the effect of chaotic orbits in steep triaxial potentials. They found that stars are supplied to the central black hole at a rate proportional to the fifth power of the stellar velocity dispersion and that the decay rate of a central black hole binary would be enhanced even if only a few percent of the stars are on chaotic orbits, thus solving the final parsec problem.  As another example, it was suggested that a third SMBH closely interacting with a hard SMBH binary can reduce the binary separation to a small value either due to the eccentricity oscillations induced in the binary via the Kozai-Lidov mechanism \\citep{2002ApJ...578..775B} or due to repopulation of the binary's loss cone due to the perturbation in the large-scale potential caused by the third black hole \\citep{2007MNRAS.377..957H}. \\citet{2002ApJ...578..775B} found that the merger time scale of an inner circular binary can be shortened by as much as an order of magnitude, and that general relativistic precession does not destroy the Kozai-Lidov effect for hierarchical triples that are compact enough.  In summary, there is substantial uncertainty in the current understanding of the evolution of binary SMBHs.  Clearly, if the SMBH binary coalescence time is longer than the typical time between successive major mergers of the galaxy, then more than two SMBHs may exist in the nucleus of a merger remnant.  We study this possibility in this paper.  We calculate the relative likelihood of binary, triple, and quadruple SMBH systems, by considering the timescales for relevant processes and combining galaxy merger trees with direct-summation N-body simulations for the dynamics of stars and SMBHs in galactic nuclei.  An obvious question regarding galactic nuclei with multiple SMBHs is whether such systems can be long-lived. We consider this question here.  Finally, systems with multiple SMBHs are likely to be interesting because of observational effects involving their effect on the properties of the host bulge, the enhancement in the rate of tidal disruption of stars, their associated gravitational wave and electromagnetic signals, and the slingshot ejection of SMBHs at high speeds.  We study some of these effects.  In \\S \\ref{prev} we review previous results on galactic nuclei with more than two SMBHs.  We present simple analytical arguments regarding the formation and evolution of such systems in \\S \\ref{sec:formation_analytic} and \\ref{sec:evolution_analytic}. Details about our numerical simulations are described in \\S \\ref{sec:simulations}, with their results shown in \\S \\ref{sec:results_numerical}.  We consider the observational signatures of our findings in \\S \\ref{sec:signatures}.  Finally, we discuss and summarise our primary findings in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we have addressed the formation of galactic nuclei with mutiple SMBHs.  We performed accurate N-body simulations of mergers of galactic nuclei with SMBHs in a cosmological setting.  Our calculation uniquely incorporated cosmological mergers of galaxies with an accurate treatment of dynamical interactions between SMBHs and stars, which we achieved using the direct summation N-body code, NBODY6.  The need for such simulations has been recognized in the literature \\citep{2005LRR.....8....8M}.  Our main conclusions are as follows:  \\begin{itemize} \\item In the absence of gas, high mass galaxies ($M_0 \\gtrsim 10^{14}$ M$_\\odot$ at $z=0$) are generically expected to have had multiple SMBHs in their nuclei during their assembly history.  Our simulations suggest that $\\sim 30$\\% galaxies within haloes with a present-day mass of $M_0\\approx 10^{14}$ M$_\\odot$ ($M_{z=6}\\approx 10^{11}$ M$_\\odot$) contain more than two SMBHs at redshifts $2\\lesssim z\\lesssim 6$.  For more massive haloes, with $M_0\\gtrsim 10^{15}$ M$_\\odot$ ($M_{z=6}\\approx 10^{12}$ M$_\\odot$), this fraction is almost 60\\%.  This is in contrast to lower-mass galaxies ($M_0 \\approx 10^{12}$ M$_\\odot$; $M_{z=6}\\approx 10^{10}$ M$_\\odot$), which rarely host more than two SMBHs in their nuclei at any moment in their assembly history.  \\item High mass galaxies as well as their low mass counterparts are rarely expected to retain more than two SMBHs in their nuclei at the present epoch.  SMBH coalescence and ejection reduces the number of SMBHs on the time scale of a Gyr.  Furthermore, major mergers are rare at lower redshift.  We also find that the number of SMBHs in galactic nuclei is rarely reduced to zero at $z=0$.  Less than 5\\% of our high-mass runs resulted in such galaxies.  \\item SMBH coalescence is common at high redshifts. Subsequent recoil due to anisotropic gravitational wave emission often results in escaping SMBHs.  Some of these SMBHs add to the wandering population of black holes in the galactic halo.  In a few cases, this process also results in galactic nuclei with no SMBH near their centres. BH-BH interaction also leads to ejected SMBHs via the slingshot mechanism.  While most of ejected SMBHs have velocities $\\lesssim 500$ km s$^{-1}$, about 10\\% SMBHs are ejected at very high velocities exceeding 2000 km s$^{-1}$.  We also find binary SMBH ejection in $\\lesssim 10 \\%$ of the cases.  \\item Multiple SMBHs have a strong effect on the stellar distribution due to three-body interactions and core passages. The resulting mass deficit is usually much larger than that due to a single SMBH binary because of resonant BH-BH interactions and GW recoil of the BH remnant.  We observe long-term oscillations of the BH-core system that could explain observations of offset AGNs.  This has implications for recent observations by \\citet{2010ApJ...717..209C} of a $z=0.359$ system that potentially contains a recoiled BH.  \\item The presence of multiple SMBHs will have important effects on the rate of tidal disruption of stars in galactic nuclei due to enhanced tidal disruption cross section, scattering of stars by other BHs, prompt loss cone refilling due to GW recoil and gravitational slingshot.  Similarly, the presence of more than two BHs in a hierarchical triple is expected to leave a signature in the GW emission from the inner binary.  This signature could be observed with future GW observatories, such as LISA.  Finally, we also expect such systems to give rise to a distinct population of wandering SMBHs that could travel in large haloes over long time scales of a few Gyrs. \\end{itemize}  The presence of gas could alter the above picture to some extent. However, simulations of binary BHs in gaseous environment have not reached sufficient resolution to confirm this.  Moreover, we expect that at high redshifts, AGN activity triggered by galaxy mergers could efficiently drive gas away from the shallow potential wells of the galaxy.  Our work can also be extended by calculating late stages of binary SMBH evolution more consistently.  New regularization techniques to do this are now available \\citep{2003gnbs.book.....A}; we defer their use to future work. Furthermore, multiple SMBH systems can also form in additional ways, for example by fragmentation of disks \\citep{2004ApJ...608..108G}.  However, these systems would evolve by migration \\citep{2011arXiv1104.2322K} on a much shorter time scale than considered here."
    },
    "1107/1107.3514_arXiv.txt": {
        "abstract": "We study the role resonant scattering plays in the transport of $\\lya$ photons in accreting protoplanetary disk systems subject to varying degrees of dust settling. While the intrinsic stellar FUV spectrum of accreting T Tauri systems may already be dominated by a strong, broad $\\lya$ line ($\\sim$80\\% of the FUV luminosity), we find that resonant scattering further enhances the $\\lya$ density in the deep molecular layers of the disk.  $\\lya$ is scattered downwards efficiently by the photodissociated atomic hydrogen layer that exists above the molecular disk.  In contrast, FUV-continuum photons pass unimpeded through the photodissociation layer, and (forward-)scatter inefficiently off dust grains.  Using detailed, adaptive grid Monte Carlo radiative transfer simulations we show that the resulting $\\lya$/FUV-continuum photon density ratio is strongly stratified; FUV-continuum dominated in the photodissociation layer, and $\\lya$-dominated field in the molecular disk. The enhancement is greatest in the interior of the disk ($r\\sim1$AU) but is also observed in the outer disk ($r\\sim100$AU). The majority of the total disk mass is shown to be increasingly $\\lya$ dominated as dust settles towards the midplane. ",
        "introduction": "}  T Tauri stars (TTS) are pre-main sequence low mass stars frequently described as young analogs of our Solar System. Surrounded by circumstellar disks of gas and dust, these systems mark the earliest stages of planet-formation, a process now understood to be common in our Galaxy \\citep{Adams:1987qf,Kenyon:1995uq,Pollack:1996fk}. Due to their close proximity to the parent star, intense ultraviolet and X-ray radiation fields greatly impact the evolution of the protoplanetary disk.  The far-ultraviolet (FUV, $h\\nu<13.6$eV) field is of particularly broad interest as it provides thermodynamical gas heating through the photoelectric effect \\citep{Weingartner:2001jl}, and affects disk chemistry directly through the photo-dissociation and -ionization of molecules and atoms \\citep[][and references therein]{Bergin:2007dz}. FUV photodesorption of ices contributes to the transport of material between solid and gas phases, whilst enabling complex molecule formation on grains surfaces \\citep{Oberg:2009kh,Oberg:2009cq}.  Alongside X-rays and cosmic-rays, the FUV field is also involved in the ionization balance of the disk; of central importance in dynamical processes such as the magneto-rotational instability that is believed to play a role in accretion \\citep{Balbus:1991fv,Gammie:1996ff}. Eventually the dispersal of disk gas may be driven by UV photo-evaporation \\citep{Alexander:2006bs}, although resolving the contributions made by the separate FUV, EUV and X-ray components is still a matter of debate \\citep{Gorti:2009ve,Owen:2010ly}   One feature of the FUV spectra in T Tauri systems which separates the accretors from the non-accretors,  is a strong, broad $\\lya$ line (accounting for as much as 80\\% of the total FUV luminosity). In nearby systems such as TW Hydra the $\\lya$ line can be detected directly \\citep{Herczeg:2002fu}, whilst in obscured systems its presence may be inferred through the excitation of $\\hh$ molecules \\citep{Herczeg:2004kb,Bergin:2004if}. It was soon realized that the concentration of UV photons in a relatively narrow wavelength range ($1215.67\\pm1$ \\AA) would have important implications for wavelength-dependent chemical processes, specifically the photodissociation of molecules \\citep{Bergin:2003la}.  For example, the molecules HCN, OH, H$_2$O and H$_2$CO can all be photodissociated by a $\\lya$ photon, unlike the closely related species CN, OH$^+$, H$_2$O$^+$ and HCO$^+$ \\citep[for a more complete list, see ][]{vDishoeck:2006kc}.  Abundance ratios frequently used as chemical diagnostics, such as $N(\\textrm{CN})/N(\\textrm{HCN})$,  will be changed (in this case increased) by $\\lya$-dominated photodissociation.  \\citet{Fogel:2011ys} have shown quantitatively that the inclusion of differential photodissociation by $\\lya$ can  some instances help bring abundance ratios into better agreement with observation \\citep{Thi:2004vn}. While the the presence of $\\lya$ decreases the abundance of  HCN, NH$_3$, and CH$_4$ by an order of magnitude or more, other species such as CO$_2$ and SO are enhanced.  SO actually has a significant cross-section at 1216$\\AA$, however so does SO$_2$, and it is the photodissociation of SO$_2$ which partially replenishes the SO abundance.  In this instance, the presence of $\\lya$ actually increases the SO abundance.  Similarly, the abundances of H$_2$O and OH in the \\citet{Fogel:2011ys} models increased with the introduction of $\\lya$, even though both H$_2$O and OH have a significant photodissociation cross-sections at 1216$\\AA$.  In this case the gas phase abundances are more than adequately replenished by efficient $\\lya$ photodesorption from grain surfaces.  Further chemical modeling is required before the implications of $\\lya$-dominated photochemistry are fully understood.  Despite the interesting chemical implications of a $\\lya$ dominated FUV field, the problem of precisely how it penetrates the disk has received little attention in the literature.  In this paper we attempt to shed light upon its basic transport.  Although frequently flared, protoplanetary disks remain relatively flat objects with vertical scaleheights that are small compared to the radial extent of the disk \\cite{Kenyon:1987kx}. As a result, the majority of the mass (residing near to the disk midplane) is concealed from the parent star. Visual optical depths at the midplane, measured along lines directly from the star, typically exceed $10^{6}$, implying that essentially no stellar photons can penetrate these parts directly. Instead, radiation must be scattered downwards from the surface of the disk \\citep{van-Zadelhoff:2003zt}. Scattering is therefore of central importance when discussing the transport of stellar radiation in protoplanetary disks.  The penetration and absorption of X-ray and FUV radiation determines much of the thermodynamical stratification of the upper disk \\citep[e.g.][and references therein]{Aresu:2011tg}. For simplicity we can identify three layers; the uppermost layer being a photodissociation layer illuminated directly by the intense, unattenuated stellar radiation field - this layer is dominated by atoms, since the photodissociation timescale of molecules is very short.  As unscattered photons penetrate into the disk they eventually reach an irradiation surface, defined to be where the optical depth integrated from the star is of order unity, $\\tau^*\\approx1$ \\citep{Calvet:1991pb,Chiang:1997kl,Watanabe:2008lo}.  It is here that, in an average sense, photons undergo their first interactions with the disk.  In the case of interactions with dust grains, this results in attenuation through both absorption and scattering; the latter leading to the emergence of a diffuse component to the radiation field. Since the radiation field decreases with depth, while the local mass density increases, we eventually enter an intermediate warm molecular layer, where $\\hh$, CO and a host of other gas-phase molecules may exist in a dynamic equilibrium with the photodissociating FUV field. It is in this warm molecular layer that the photochemical processes are most important. Despite its attenuation, the radiation field is sufficient to maintain grain temperatures above $30$K, thus preventing widespread freeze-out of molecules onto grain surfaces.  At even greater depths we enter the optically thick, dense midplane region containing $>99\\%$ of the disk mass.  Here the stellar FUV field is highly attenuated and consists entirely of scattered photons.  With the exception of the innermost few AU, the temperatures may be low enough ($<30$K) to permit the freeze-out of several molecular species \\citep{Aikawa:2002hs,Bergin:2007dz}.  The transport of FUV photons is further complicated by the gradual growth and settling of dust grains as the disk evolves towards planetesimal formation \\citep{Throop:2001ee,Wilner:2005kn}.   Not only does grain-growth affect the optical properties of individual grains, but the radial and vertical drift of solids relative to gas causes a segregation of the grain-size population \\citep{Weidenschilling:1977xw,Dullemond:2004mi}.  The majority of the grain opacity at FUV wavelengths is attributable to relatively small grains with radii $a<1\\mu$m, which subsequently reemit absorbed energy at infrared wavelengths, making it possible to estimate their abundance \\citep{Dominik:2008hc}.  Observationally, the abundance of small grains (per H nucleus) relative to that found in the ISM, $\\epsilon$, typically takes a value much less than one; for example, in the Taurus star-forming region the median value is $\\epsilon\\sim0.01$ \\citep{Furlan:2006fe}.  %  Our basic picture of FUV transport is somewhat incomplete once we include $\\lya$ photons.  In view of the large cross-section due to resonant scattering by H atoms, $\\sigma_{\\lya}$, it seems entirely plausible that $\\lya$ photons experience their first interactions not with dust grains but with H atoms residing in the uppermost photodissociation layer of the disk (hereafter `H-layer').  The irradiation surface for $\\lya$ photons is therefore defined as the location where  $\\tau^*=\\int \\sigma_{\\lya} n(H)\\approx1$.  The precise location depends quite sensitively on the width of the stellar $\\lya$ line profile, although the $\\lya$ photons that propagate freely to the disk surface are generally those in the wings of the scattering line profile presented by the disk \\citep{Herczeg:2002fu}.   Even hundreds of Doppler widths out in the Lorentz wings of the scattering line profile the cross-section may still be significantly larger than that due to the (possibly settled) dust population.  Figure \\ref{fig:voigt} shows the $\\lya$ line from TW Hya, plotted next to an estimate of the scattering line profile due to warm H atoms in the disk.   Assuming the photodissociation layer supports H columns of $N^*(H)\\ge10^{20}$cm$^{-2}$ (integrated along a line from the star), it seems probable that the $\\lya$ field will first be scattered by this atomic H-layer.  The optical depth of dust in this layer will be smaller by a factor of approximately $0.1\\epsilon<<1$.  %  In stark contrast to $\\lya$ transport, FUV-continuum photons will stream unimpeded through the H-layer, eventually striking the disk at the conventional irradiation surface defined by dust.  Here they will begin to scatter, however scattering by typical interstellar dust populations is considered to be somewhat forward-throwing \\citep{Gordon:2004lh}.  Multiple scatterings are therefore required before an appreciably diffuse component is generated, and accompanying this scattering will be a degree of pure absorption.  %    In this paper we explore numerically the basic phenomenology of resonantly scattered $\\lya$ in disk models. An emphasis is placed on contrasting the transport of $\\lya$ with that of the FUV-continuum field in the vicinity of the $\\lya$ line.  We proceed with Section \\ref{sec:disks} where we discuss the properties of our disk models.  Radiation transfer of $\\lya$ and FUV-continuum photons is dealt with in Section \\ref{sec:Radiative-transfer}.  This includes an approximate - but largely self-consistent -computation of the H distribution necessary for $\\lya$ transport. Results are presented in Section \\ref{sec:Results}.  The paper concludes with a discussion and summary. ",
        "conclusions": "}  Regardless of the value of $\\epsilon$, essentially every stellar $\\lya$ photon is intercepted by the highly flared H-layer. Upon the first scattering the $\\lya$ field is completely isotropized and a fraction $(<50\\%)$is transmitted through to the base H-layer. This is clearly seen in Fig. \\ref{fig:maps} (top left and bottom left panels).  Viewed from within the molecular region the $\\lya$ would seem to originate in a diffuse blanket overlying the disk, analogous to the diffuse transport of sunlight on an overcast day.  This is similar to the penetrative advantage that interstellar photons (bathing the disk isotropically) would have \\citep[e.g.][]{Willacy:2000zr}, however in our models the stellar $\\lya$ field is typically orders of magnitude more intense.  Once in the realm of the molecular disk, $\\lya$ transport will be controlled primarily by dust.  FUV-continuum photons pass unimpeded through the H-layer until they reach the dust irradiation surface. Until this point the field only experiences geometric inverse-square law diminution - in contrast to the $\\lya$, which has been additionally processed by the H-layer. In this regime the $\\lya$/FUV-continuum photon density ratio favors the unimpeded FUV-continuum photons. The FUV-continuum eventually impinges upon the dust irradiation surface, doing so at very small angles ($\\leq0.05$ degrees).  The relative inefficacy of scattering by dust grains is exacerbated by accompanying absorption. It is straightforward to show with 1D planar radiative transfer models that the energy density transmitted downwards from this dust irradiation surface is typically of order $\\leq1\\%$ of that incident upon it \\citep{Chandrasekhar:1960kl}. Consequently, the $\\lya$/FUV-continuum photon density ratio undergoes a dramatic reversal in the vicinity of the dust irradiation surface, resulting in a $\\lya$-dominated field.  This resurgence appears to more than compensate for the reduction of $\\lya$ density due to the H-layer.  In the asymptotic limit (large optical depth) of the analytic planar model both fields are attenuated according to $n_{ph}\\propto\\exp\\left[-k\\tau_{z}\\right],$where $\\tau_{z}$ is the vertical optical depth due to dust, and $k$ is a decay constant that depends only on the scattering properties $(\\omega,g)$ of dust \\citep{Flannery:1980xr}.  As observed, the $\\lya$/FUV-continuum ratio approaches a constant value deep inside the disk (the `attenuated region' in Figure \\ref{fig:ratios}).  At this point both $\\lya$ and FUV-continuum fields are largely isotropized and their transport controlled by dust alone.  In effect, the contrasting early life-histories are imprinted into boundary conditions that determine the proportionality coefficients in the asymptotic regime.  A simplified diagrammatical representation of the transport routes is shown in Figure \\ref{fig:nh_map}.  As mentioned previously, the D'Alessio disks to not include separate thermodynamic treatments for gas and dust.  In effect they assume that the two components are thermally coupled.  Models show that this coupling breaks down at low densities, causing the gas temperature to rise above that of dust \\citep{Glassgold:2004ec}.  In turn this increases the local scale height of the gas while decreasing its density.    Although not explored in this paper, the thermal decoupling will most likely increase the extent of atomic H in the upper layers of the disk.  This is due to two factors; first, the formation rate of H$_2$ is density dependent and an increase in the disk scaleheight will be accompanied by a decrease in volumetric density.  Second,  H$_2$ formation on grains becomes increasingly inefficient at $T\\ge 500K$ due to the short time that H atoms spend on grain surfaces before thermally escaping \\citep{Cazaux:2004fu}.    The settling of dust affects the $\\lya$ and FUV-continuum fields in two distinct ways. First, the removal of dust decreases the efficacy of the $\\hh$ formation process, resulting in a thicker H-layer (Eqn. \\ref{eq:H2}, Fig. \\ref{fig:columns}). This ensures that the H-layer processes every stellar $\\lya$ incident upon it.  Second, the irradiation surface for FUV-continuum photons recedes into the disk as $\\epsilon$ decreases.  The disk becomes flatter and intercepts a smaller fraction of stellar FUV-continuum photons.  A self-consistent disk model that couples disk structure with our new radiative transfer results is required in order to quantify these effects more precisely.  We conclude with a brief summary of the main results: \\begin{enumerate} \\item The vertical profile of the $\\lya$/FUV-continuum photon density ratio is strongly stratified, and may vary by orders of magnitude relative to the intrinsic stellar ratio. The stratification is strongest in the inner disk. \\item The radiation field in the vicinity of the disk midplane - where most of the disk mass resides - is relatively enhanced in $\\lya$.  This enhancement is in addition to that already present in the stellar spectrum. \\item If previous radiative transfer calculations most closely resemble our FUV-continuum results, then the omission of resonantly scattered $\\lya$ implies that the mass-averaged intradisk FUV field has been systematically underestimated.  This directly affects estimates of photoelectric heating rates, thermal balance, and the resulting hydrostatic structure. \\item The settling of dust ($\\epsilon<1$) lowers the dust irradiation surface and thickens the resonant scattering H-layer by inhibiting $\\hh$ formation.  Both effects contribute to a more stratified  $\\lya$/FUV-continuum ratio, ultimately resulting in a greater $\\lya$ enhancement in the molecular disk. \\item The presence of an optically thick, high-albedo H-layer implies that the direct observation of $\\lya$ scattered by the disk ($\\geq50\\%$ of incident photons) may be possible in sufficiently close, low extinction systems (e.g. TW Hya). These observations may place constraints on the shape and thickness of the H-layer. \\item A $\\lya$-dominated radiation field drives differential photo-chemical processes in both the gas and grain surfaces \\citep{Bergin:2003la,Fogel:2011ys}. Numerous pairs of closely-related molecules exhibit a similar differential response to the spectral form of the UV field \\citep{vDishoeck:2006kc}. In view of the strongly stratified $\\lya$/FUV-continuum ratio one might also expect the photochemical environment to be similarly stratified. \\end{enumerate}"
    },
    "1107/1107.3949_arXiv.txt": {
        "abstract": "The observation of \\emph{blank fields}, regions of the sky devoid of stars down to a given threshold magnitude, constitutes one of the typical important calibration procedures required for the proper reduction of astronomical data obtained in imaging mode. This work describes a method, based on the use of the Delaunay triangulation on the surface of a sphere, that allows the easy generation of blank fields catalogues. In addition to that, a new tool named TESELA, accessible through the WEB, has been created to facilitate the user to retrieve, and visualise using the VO-tool Aladin, the blank fields available near a given position in the sky. ",
        "introduction": "\\label{section:introduction}  Scientific research progress is critically based on the correct exploitation of data obtained at the limit of the technological capabilities.  Thus prior to data analysis and interpretation, the quality of the data treatment must be assured in order to guarantee that the information content in that data is preserved. In observational Astrophysics, and in particular imaging with ground-based telescopes, data treatment is performed following data reduction procedures. The main goal of such reduction processes is to minimize the influence of data acquisition imperfections on the estimation of the desired astronomical measurements (see e.g.\\ \\citealt{gilliland92} for a short review on noise sources and reduction processes of CCD data). For this purpose, one must perform appropriate manipulations of data and calibration frames, the latter properly planned and obtained to facilitate the reduction procedure. In this sense, image flatfielding and sky subtraction constitute two of the most common and important reduction steps.  Their relevance should not be underestimated, since inadequate flatfielding or sky subtraction easily leads to the introduction of systematic uncertainties in the data. Contrary to random errors, which can be controlled by applying normal statistical methods, the systematic uncertainties can propagate hidden within the arithmetically manipulated data to the final measurements, where their impact may considerably bias their interpretation.  Concerning image flatfielding, the corresponding calibration frames needed to proceed through this reduction step are typically divided into two categories: images intended to provide high-frequency scale sensitivity variations (the pixel-to-pixel response), and images whose goal is to facilitate the removal of the two-dimensional low-frequency scale sensitivity variations of the detectors. The former are usually obtained from continuum lamps (either within the instrument itself or by illuminating a flat screen or the inner dome of the telescope), although the colour of such lamps are not expected to match the spectral energy distributions of the science targets, and ideally dark night sky flats, obtained by pointing the telescope to a \\emph{blank field} should be observed instead. Since this is quite expensive in terms of observing time, this approach is very rarely used. The low-frequency flatfields are usually twilight exposures. The problem here is that the time during which the sky surface brightness is high enough to clearly overpass the signal of any star in the field of view, but not too high to saturate the frames, are placed on two time windows every day, one at the beginning of the evening twilight and the other at the end of the morning twilight. The observational strategy typically consist on obtaining a series of images, shifting the telescope position a few arcseconds between the different exposures in order to avoid the bright stars in the field of view to appear in the same detector pixels. Those pixels are afterwards masked when combining the individual exposures during the reduction procedure.  In any case very bright stars need to be avoided since the point spread function of these objects can introduce spikes and artifacts that are not so easily removed during the reduction of the flatfield images. For this purpose the use of \\emph{blank fields} is the best option.  Sky subtraction cannot always be performed by measuring the night sky level in the same frame where the scientific targets are present. This is especially true when the target dimensions are comparable to the detector field of view. In that circumstance separate sky frames are observed, by pointing the telescope to a position devoid of stars as much as possible.  For the sky subtraction procedure to be successful, it is essential that these separate sky images are obtained under identical observing conditions, i.e.\\ very close in time and sky position, since this is the only way to prevent varying observing conditions that modify the sky surface brightness.  From the above discussion it is clear that the observation of sky regions devoid of stars down to a given magnitude is a very important aspect to be taken into account when carrying out astronomical observations. To date, no systematic catalogue is available providing the location in the sky of blank fields for varying limiting magnitudes\\footnote{One of the scarce resources is the web page created by Marco Azzaro at {\\tt http://www.ing.iac.es/\\~{}meteodat/blanks.htm}, where a list of 38 blank fields is available.}. With modern instrumentation intended to be used in large imaging surveys, the problem is expected to become more demanding. The situation is also starting to be important in spectroscopic observations, since the advent of integral field units with high multiplexing capabilities (and thus, increasingly larger field of view), are expected to be one of the most common instrumentation in ground-based observatories.  The work presented in this paper describes a method that helps to determine the availability of blank fields in any region of the celestial sphere, based on the use of the Delaunay triangulation. The method has been employed to generate catalogues of blank fields to varying limiting magnitudes. In order to facilitate the use of these catalogues, we have created TESELA, a new tool accessible through the WEB which provides a simple interface that allows the user to retrieve the list of blank fields available near a given position in the sky.  Section~\\ref{section:tessellating-the-sky} of this paper describes the method and presents some statistical analysis of its application to the \\mbox{Tycho-2} \\citep{Hog00a} catalogue. Section~\\ref{section:TESELA} describes the new tool TESELA. The detailed description of the tessellation of sky subregions is presented in the Appendix. ",
        "conclusions": ""
    },
    "1107/1107.4272_arXiv.txt": {
        "abstract": "During its first year of data taking, the Large Area Telescope (LAT) onboard the {\\em Fermi} Gamma-Ray Space Telescope has collected a large sample of high-energy cosmic-ray electrons and positrons (CREs). We present the results of a directional analysis of the CRE events, in which we searched for a flux excess correlated with the direction of the Sun. Two different and complementary analysis approaches were implemented, and neither yielded evidence of a significant CRE flux excess from the Sun.  We derive upper limits on the CRE flux from the Sun's direction, and use these bounds to constrain two classes of dark matter models which predict a solar CRE flux: (1) models in which dark matter annihilates to CREs via a light intermediate state, and (2) inelastic dark matter models in which dark matter annihilates to CREs. ",
        "introduction": "In the last decades the searches for a dark matter (DM) signal from the Sun were performed looking for possible excesses of neutrinos or gamma-rays associated with the Sun's direction. However, as it was noted in Ref.~\\cite{Schuster:2009fc}, several DM models that have been recently developed to explain various experimental results also imply an associated solar flux of high-energy cosmic-ray electrons and positrons (CREs). On the other hand, no known astrophysical mechanisms are expected to generate a significant high-energy CRE $(>100 \\units{GeV})$ excess associated with the Sun.  A class of models in which DM annihilates to CREs through a new light intermediate state $\\phi$ \\cite{Pospelov:2007mp,ArkaniHamed:2008qn} has been considered to explain the excesses in local CRE fluxes reported by PAMELA~\\cite{Adriani:2008zr}, ATIC~\\cite{:2008zzr}, and \\emph{Fermi}~\\cite{Abdo:2009zk,Ackermann:2010ij}. In these scenarios DM particles captured by the Sun through elastic scattering interactions would annihilate to $\\phi$ pairs in the Sun's core, and if the $\\phi$ could escape the surface of the Sun before decaying to CREs, these models could produce an observable CRE flux.  Another class of models in which DM scatters off of nucleons predominantly via inelastic scattering has been proposed as a means of reconciling the results of DAMA and DAMA/LIBRA~\\cite{Bernabei:2008yi,Bernabei:2010mq} with CDMS-II~\\cite{Ahmed:2009zw,Ahmed:2010hw} and other experiments (e.g., \\cite{Chang:2008gd,Finkbeiner:2009ug}; see also \\cite{Savage:2008er} for a comprehensive discussion of experimental constraints).  If DM is captured by the Sun only through inelastic scattering (iDM), this could lead to a non-negligible fraction of DM annihilating outside of the Sun's surface.  For models in which iDM annihilates to CREs, an observable flux at energies above a few tens of GeV could be produced.  During its first year of operation, the Large Area Telescope (LAT) onboard the {\\em Fermi} satellite~\\cite{Atwood:2009ez} has collected a substantial number of CRE events, which has allowed a precise measurement of the energy spectrum over a broad energy range from a few $\\units{GeV}$ up to $1\\units{TeV}$~\\cite{Abdo:2009zk,Ackermann:2010ij}. Furthermore, a directional analysis of the high-energy CRE events was performed in the Galactic reference frame~\\cite{Ackermann:2010ip}, and showed no evidence of anisotropies.  In this paper we use the high-energy CRE data set to search for flux variations correlated with the Sun's direction. Since the Sun is moving with respect to the Galactic reference frame, the previously-reported absence of anisotropies in the CRE flux observed in the Galactic frame does not necessarily imply a negative result. ",
        "conclusions": "We used a sample of about $1.3 \\times 10^{6}$ CRE events with energies above $60\\units{GeV}$ detected by the {\\em Fermi} LAT during its first year of data-taking to search for  flux excesses or deficits correlated with the Sun's direction. Two analysis approaches were implemented, and neither yielded evidence of an enhancement in the CRE flux from the direction of the Sun. This result agrees with the more general one shown in Ref.~\\cite{Ackermann:2010ip}, where no evidence of anisotropies was found in CRE arrival directions above $60\\units{GeV}$ in the Galactic reference frame.  We derived limits on DM models which generate a CRE flux from the Sun's direction for the two scenarios discussed in Ref.~\\cite{Schuster:2009fc}. In the case of annihilation of DM through an intermediate state and subsequent decay to $e^{\\pm}$, the upper limits on solar CRE fluxes provide significantly stronger constraints on the DM scattering cross-section than limits previously derived by constraining the FSR emission associated with this decay channel using solar gamma-ray measurements. For the iDM scenario, the solar CRE flux upper limits exclude the range of models which can reconcile the data from DAMA/LIBRA and CDMS for $m_{\\chi} \\gtrsim 70\\units{GeV}$, assuming DM annihilates predominantly to $e^{\\pm}$.  Since direct detection experiments are not sensitive to the dominant annihilation channels of the DM particles, other data, e.g., solar gamma-ray measurements and neutrino searches, may be able to further constrain these models by excluding regions of parameter space for alternative annihilation channels."
    },
    "1107/1107.2498.txt": {
        "abstract": "We report on near-infrared $J$- and $H$-band linear polarimetric photometry of eight ultracool dwarfs (two late-M, five L0--L7.5, and one T2.5) with known evidence for photometric variability due to dust clouds, anomalous red infrared colors, or low-gravity atmospheres. The polarimetric data were acquired with the LIRIS instrument on the William Herschel Telescope. We also provide mid-infrared photometry in the interval 3.4--24 $\\micron$ for some targets obtained with {\\sl Spitzer} and {\\sl WISE}, which has allowed us to confirm the peculiar red colors of five sources in the sample. We can impose modest upper limits of 0.9\\,\\%~and 1.8\\,\\%~on the linear polarization degree for seven targets with a confidence of 99\\,\\%. Only one source, 2MASS\\,J02411151$-$0326587 (L0), appears to be strongly polarized ($P \\sim 3\\,\\%$) in the $J$-band with a significance level of $P/\\sigma_P \\sim 10$. The likely origin of its linearly polarized light and rather red infrared colors may reside in a surrounding disk with an asymmetric distribution of grains. Given its proximity (66\\,$\\pm$\\,8 pc), this object becomes an excellent target for the direct detection of the disk. ",
        "introduction": "There is growing evidence that dwarfs of spectral types L and T show a large spread of near- and mid-infrared colors that cannot be explained by simple models (e.g., Marley et al$.$ \\cite{marley10}). It is believed that gravity, metallicity, and distribution of dust clouds play a critical role in defining the atmospheric properties of L- and T-type sources (see Kirkpatrick \\cite{kirk05} for a review). Various groups have monitored photometrically a number of L and T objects aimed at characterizing cloud patchiness in brown dwarfs (e.g., Gelino et al$.$ \\cite{gelino02};  Koen \\cite{koen04a}; Koen et al$.$ \\cite{koen04b}; Artigau et al. \\cite{artigau06}). Polarimetric observations at optical wavelengths have also been attempted to confirm the presence of dusty clouds in the objects' atmospheres (M\\'enard et al$.$ \\cite{menard02}: Zapatero Osorio et al$.$ \\cite{osorio06}; Goldman et al$.$ \\cite{goldman09}; Tata et al$.$ \\cite{tata09}). The main results are that photometric variability (if detected) has a relatively small amplitude (typically below 0.1 mag in the $I$-band and infrared wavelengths), variability can be periodic (with modulations in agreement and disagreement with the expected rotation periods), non-periodic variability also appears quite commonly, and a few percentage of the L dwarfs show linear polarization degrees typically below 1\\%~in the optical bands. All these observations support the presence of condensates in the form of dust clouds and the fast rotation of the dwarfs.  Among the known L dwarfs, there are some showing markedly redder near- and mid-infrared colors than expected for their spectral type . The origin of this property is not well understood. Burrows et al$.$ \\cite{burrows06}, Leggett et al$.$ \\cite{leggett07}, Burgasser et al$.$ \\cite{burgasser08}, and Cruz et al$.$ \\cite{cruz09} argued that low gravity causes red colors, but not all red L dwarfs necessarily have low surface gravities. An excess of metallicity and/or unusual condensate properties can also be the cause of a reddish nature. In Zapatero Osorio et al$.$ \\cite{osorio10} we explored various scenarios for one particular very red L dwarf: G\\,196$-$3\\,B. We concluded that a low-gravity atmosphere with enshrouded upper atmospheric layers and/or a warm dusty disk/envelope provides the most likely explanations, the two of them consistent with an age below 300 Myr.  Imaging polarimetry provides a powerful tool for studying the properties of the atmospheric dust clouds and surrounding disks/envelopes. While the models by Sengupta \\& Marley \\cite{sengupta09,sengupta10} (and references therein) predict non-zero disk-integrated polarization with degrees quite below 1\\%~for the L and T dwarfs, polarimetric observations of young stars with disks yield higher linear polarization intensities (e.g., Kusakabe et al$.$ \\cite{kusakabe08}; Pereyra et al$.$ \\cite{pereyra09}, and references therein). Here, we report on the first results of our polarimetric observations of ultracool dwarfs (spectral types later than M7) with peculiar colors and photometric variability or low gravity atmospheres. ",
        "conclusions": "We carried out polarimetric imaging observations of eight ultracool dwarfs (most likely brown dwarfs) with spectral types in the interval M7--T2.5 using the $J$- and $H$-band filters and the LIRIS instrument on the WHT. The target list is formed by objects with detected photometric variability due to dust clouds, very red $J-K_s$ colors (as compared to other field dwarfs of similar classification), and low gravity atmospheres. We also obtained {\\sl Spizter} IRAC and MIPS photometry (from the public archive) for five objects in the sample; in addition, four sources are included in the recent release of the {\\sl WISE} catalog. This has allowed us to study the $J-[4.5]$ color of all targets in comparison to the mean color of field dwarfs with related spectral types, finding that five sources show clear $J-[4.5]$ color excesses, four of which are spectroscopically confirmed young objects with likely ages below 300 Myr. From the LIRIS data, we can impose 99\\%-confidence level upper limits of $P \\le 0.9$\\%~on the near-infrared linear polarization degree for J2244$+$20, J0136$+$09, and J1022$+$58, and $P \\le 1.8$\\%~for G\\,196$-$3\\,B, J0355$+$11, and USco\\,128 and USco\\,132. We detected significant $J$-band linear polarization ($P = 3.04\\pm0.30$) in the L0 dwarf J0241$-$03, which is one of the young brown dwarfs with notorious infrared color excesses: $E(J-[4.5]) = 0.73\\pm0.23$ mag. The likely scenario to account for both the infrared flux excesses and the polarized light is the presence of a surrounding dusty envelope or disk with an asymmetric distribution of grains. The study of the linear polarization degree and polarization angle at different wavelengths may shed new light on the properties of the grains responsible for the observed polarization in J0241$-$03. The incidence of positive linear polarimetric detections with degree $P \\ge 1.8$\\%~in our sample of very red young sources ($\\le$300 Myr)  is $25\\pm25$\\%."
    },
    "1107/1107.3859_arXiv.txt": {
        "abstract": "Pions may remain stable under certain conditions in a dense media at zero temperature in the normal phase (non pion superfluid state). The stability condition is achieved when the in-media pion width vanishes. However, thermal fluctuations will change this stable regime. For low temperature  pions will remain in a metastable state. Here we discuss the different possible scenarios for leptonic pion decays at finite temperature, taking into account all the different chemical potentials involved. The neutrino emission due to pions in a hot-dense media is calculated, as well as the coolig rate of a pion-lepton gas. ",
        "introduction": "\\label{LEQCD}  As the 99.9877\\% of the charged pions decay into muons and muonic neutrinos, we will refer specifically to this process. Our discussion will be based on the following low-energy Lagrangian: ${\\cal L}= {\\cal L}_{\\mu}+ {\\cal L}_\\chi$, where \\begin{equation} {\\cal L}_{\\mu} = \\sum_{f=\\mu,\\nu_\\mu} \\bar \\psi_f \\left(i\\slashed{\\partial}+\\mu_f\\gamma_0 +m_f\\right)\\psi_f \\end{equation} corresponds to the lepton free Lagrangian, $\\mu_f$ being the associated lepton chemical potential. The second term in the low-energy Lagrangian corresponds to the ${\\cal O}(p)^2$ chiral Lagrangian \\cite{Gasser:1983yg, Donoghue:1992dd} \\begin{equation} {\\cal L}_{\\chi} = \\frac{F^2}{4}\\textrm{tr}(D_\\mu U)^\\dag DU^\\mu +\\frac{G}{2} \\textrm{tr}(U^\\dag{\\cal M}+{\\cal M}^\\dag U). \\label{Lchi} \\end{equation} The $F$ and $G$ terms are the tree level pion decay constant and the tree level chiral condensate, respectively. The $U$ fields contain the pion fields as $U=\\exp(i\\pi^a\\tau^a/f)$ and the covariant derivative $D_\\mu U=\\partial_\\mu U-ir_\\mu U+iUl_\\mu$ includes external right and left currents. When ${\\cal M}=\\textrm{diag}(m_u,m_d)$ we break explicitly the chiral symmetry. In this article we are not interested in computing pion radiative corrections. Therefore, higher order chiral Lagrangian terms will not be considered. However, those corrections have been already calculated at finite temperature  and isospin chemical potential  \\cite{Loewe:2002tw}, and therefore, if we want to incorporate these contributions, it is enough to replace the masses, decay constant and  chemical potential terms by temperature and isospin chemical potential dependent dressed terms  \\cite{Fraga:2008be}.  We will use the effective Fermi model for the leptonic weak coupling. The leptonic weak currents and the isospin chemical potential are introduced by setting the external currents in the chiral Lagrangian as \\begin{eqnarray} r_\\alpha &=& \\frac{1}{2}\\mu_I\\tau^3\\delta_{\\alpha 0}\\\\ l_\\alpha &=& G_F[\\bar\\psi_{\\nu_\\mu}\\gamma_\\alpha(1-\\gamma_5)\\psi_{\\mu}~\\tau^- +\\bar\\psi_{\\mu}\\gamma_\\alpha(1-\\gamma_5)\\psi_{\\nu_\\mu}~\\tau^+] \\nonumber\\\\&& +\\frac{1}{2}\\mu_I\\tau^3 \\delta_{\\alpha 0} \\end{eqnarray} with $\\mu_I=\\mu_u-\\mu_d$ is the isospin chemical potential and where we use the following combination of Pauli matrices: $\\tau^\\pm =\\frac{1}{\\sqrt{2}}(\\tau^1\\pm i\\tau^2)$. We will concentrate only on the normal phase where $|\\mu_I|< m_\\pi$. In the superfluid phase, $|\\mu_I|> m_\\pi$, one of the charged pions condenses, and therefore, another treatment is needed \\cite{Son:2000xc,Loewe:2004mu}.   \\subsection*{Baryon chemical potential}  Two different regions in the literature have been considered when the Baryon chemical potential $\\mu_B=\\frac{3}{2}(\\mu_u+\\mu_d)$ is introduced in the frame of chiral perturbation theory:  \\emph{Small} $\\mu_B$. The baryon chemical potential can be taken as ${\\cal O}(p)$ in the power counting in chiral perturbation theory. In this case $\\mu_B$ appears in the Wess-Zumino-Witten anomalous term, which turns out to be relevant only in higher order radiative corrections \\cite{AlvarezEstrada:1995mh}.  \\emph{Very high} $\\mu_B$. In this case an expansion in powers of ${\\mu_B}^{-1}$ is performed (asymptotically infinite chemical potential). Here, effects of color superconductivity and color flavor locked phase are present, due to the appearance of diquark pairing \\cite{Hong:1999dk,Casalbuoni:1999wu,Son:1999cm,Beane:2000ms}. We can construct an effective Lagrangian including $\\mu_B$ in the absence of diquark effects, by considering effective $F(\\mu_B)$ and $G(\\mu_B)$ constants in the chiral Lagrangian in Eq.~(\\ref{Lchi}) and introducing a Lorentz symmetry breaking term: $(D_\\mu U)^\\dag DU^\\mu\\to (D_0 U)^\\dag D_0U-v^2(\\bm{D} U)^\\dag\\cdot \\bm{D}U$. Results obtained in the frame of the Nambu--Jona-Lasinio model show that for $\\mu_B\\lesssim1$~GeV the $F$, $G$ and $v$ parameters  do not suffer significant changes \\cite{Bender:1997jf, Ebert:2005wr, Jiang:2008rb}. We can neglect then baryon chemical potential effects in this work.   Once the chiral Lagrangian has been expanded in terms of the pion fields, the canonical quantization procedure is the standard one, keeping in mind that the energy of the charged pions and lepton fields are shifted due to the chemical potentials. Then, we proceed to calculate the decay width of the charged pions and the corresponding neutrino emissivity. ",
        "conclusions": "\\label{conclusions}  In this paper we have calculated the charged pions decay widths in dense matter at finite temperature, analyzing the pion metastable condition. We have calculated also the neutrino emissivity through the leptonic pion decay in $\\beta$-equilibrium for a degenerated system. We obtained the main contributions in the low temperature region for pions decaying into muons and muonic neutrinos. Finally, we estimated the cooling rate of a pion lepton gas in $\\beta$-equilibrium for three different values of the electric chemical potential.    From our results, we argue that it is possible to find stable pions, even if they are not condensed, if the lepton chemical potential reaches a value higher than the muon mass, which is the case in compact stars. Under such conditions, the pions might only decay through thermal fluctuations. The muonic neutrino emissivity will grow with the lepton chemical potential, varying from  $\\sim T^{3/2}$ to $\\sim T$, both with an exponential suppressing term. The contribution to the cooling process in a neutron star, due to the emission of  muonic neutrinos from pions, has not been much considered yet in the normal phase. In fact, from our estimation of the cooling time of the lepton pion gas, it can be seen that this process is relevant for temperatures higher than ~1 MeV which corresponds to the cooling process of a protoneutron star. The pion superfluid case will be explored elsewhere."
    },
    "1107/1107.1771_arXiv.txt": {
        "abstract": "{It is well established that the winds of carbon-rich AGB stars (carbon stars) can be driven by radiation pressure on grains of amorphous carbon and collisional transfer of momentum to the gas. This has been demonstrated convincingly by different numerical wind models that include time-dependent dust formation. To simplify the treatment of dust opacities, radiative cross sections are usually computed using the assumption that the dust grains are small compared to wavelengths around the stellar flux maximum. Considering the typical grain sizes that result from these models, however, the applicability of this small-particle limit (SPL) seems questionable.} {We explore grain size effects on wind properties of carbon stars, using a generalized description of radiative cross sections valid for particles of arbitrary sizes. The purpose of the study is to investigate under which circumstances the SPL may give acceptable results, and to quantify the possible errors that may occur when the SPL does not hold.} {The time-dependent description of grain growth in our detailed radiation-hydrodynamical models gives information about dust particle radii in every layer at every instant of time. Theses grain radii are used for computing opacities and determining the radiative acceleration of the dust-gas mixture. From the large number of models presented in the first paper of this series (based on SPL dust opacities; Mattsson et al.~2010) we selected two samples, i.e., a group of models with strong, well-developed outflows that are probably representative of the majority of wind-forming models, and another group, close to thresholds in stellar parameter space for dust-driven winds, which are referred to as critical cases.} {We show that in the critical cases the effect of the generalized description of dust opacities can be significant, resulting in more intense mass loss and higher wind velocities compared to models using SPL opacities. For well-developed winds, however, grain size effects on mass-loss rates and wind velocities are found to be small. Both groups of models tend towards lower degrees of dust condensation compared to corresponding SPL models, owing to a self-regulating feedback between grain growth and radiative acceleration. Consequently, the \"dust-loss rates\" are lower in the models with the generalized treatment of grain opacities.} {We conclude that our previous results on mass-loss rates obtained with SPL opacities are reliable within a wide region of stellar parameter space, except for critical cases close to thresholds of dust-driven outflows where SPL models will tend to underestimate the mass loss rates and wind velocities. } ",
        "introduction": "Winds of carbon stars are usually considered to {be dust-driven winds}. Stellar photons, incident on dust particles, will lead to a radiative acceleration of the grains away from the star, and, subsequently, momentum will be transferred to the surrounding gas by gas--grain collisions. Pulsation-induced atmospheric shock waves contribute significantly to this process by intermittently creating cool, dense layers of gas well above the photosphere where dust grains can form and grow efficiently.  Pioneering work on the modelling of AGB star winds was done by Wood (1979), \\nocite{Wood79} focusing on the effects of shock waves, and later by Bowen (1988), \\nocite{Bowen88} introducing a parameterized opacity to describe the dynamical effects of dust formation in the circumstellar envelope. These early wind models where followed by studies of carbon stars including time-dependent {(non-equilibrium)} grain growth (e.g. Fleischer et al 1992; H\\\"ofner \\& Dorfi 1997; Winters et al. 2000) which, despite being based on grey radiative transfer, allowed to describe basic properties of heavily dust-enshrouded carbon stars. In order to obtain reasonably realistic results for objects with less optically thick envelopes, however, it is necessary to combine frequency-dependent radiative transfer (including gas and dust opacities) with time-dependent hydrodynamics and non-equilibrium dust formation (cf. H\\\"ofner et al.~2003).  At a point where models of carbon stars are becoming quantitatively comparable to observations as diverse as high-resolution IR spectra (e.g. Nowotny et al. 2010) and spectro-interferometric measurements (e.g. Sacuto et al. 2011), it is necessary to scrutinize a number of underlying physical assumptions and approximations. In particular, this concerns the detailed treatment of dust opacities that are at the core of the wind mechanism and also have a direct influence on observable properties. A common feature of most detailed dust-driven wind models in the literature (including the first paper in this series by Mattsson et al. 2010) is that dust opacities are computed under the assumption that the grain sizes are small compared to the relevant wavelengths (defined by the stellar flux distribution), using the small-particle limit (SPL) of the Mie theory. In this limit, dust opacities are fully determined by the amount of condensed material, irrespective of grain sizes, which greatly simplifies the modelling because an explicit knowledge of the actual grain size distribution in each layer is not required. However, it has been shown that grains may grow to sizes where the use of the SPL is questionable (e.g., Gail \\& Sedlmayr 1987, Winters et al. 1994, 1997, Mattsson et al.~2010).  The ongoing debate on the mass-loss mechanism of M-type AGB stars has recently put the effects of grain size in these objects into focus. Using detailed non-grey models, Woitke (2006) demonstrated that silicate grains have to be virtually Fe-free in the wind acceleration zone, which leads to insufficient radiative pressure caused by absorption. Models by H{\\\"o}fner (2008) suggest scattering as a possible solution: if conditions in the extended atmosphere allow these grains to grow into the size range of about $0.1 - 1 \\mu$m, scattering becomes dominant over absorption by several orders of magnitude, opening up the possibility of stellar winds driven by scattering on virtually Fe-free silicate grains.  \\begin{figure} \\includegraphics[width=9cm]{aa15572_fig1} \\caption{  \\label{kappa_cont} Relevant dust opacity for radiation pressure {(combining effects of absorption and scattering, see Sect.~\\ref{s_method})} as function of grain radius and wavelength, {computed from refractive index data for amorphous carbon} by Rouleau \\& Martin (1991). The black contour shows the region where the flux-weighted monochromatic opacity exceeds the critical opacity, that is {\\bf required in order for the radiation pressure to balance gravity} (see Eqs.~\\ref{e_kappa_weight} and \\ref{crit}, respectively), assuming a Planckian flux distribution with $T_{\\rm eff} = 2700$~K and that 30\\% of the carbon not bound in CO condense into carbon dust. } \\end{figure}  In carbon stars, on the other hand, the effects of grain size are expected to be less dramatic owing to the high absorption cross sections of amorphous carbon grains. As shown in Fig.~\\ref{kappa_cont}, a wide range of particle sizes can contribute to driving winds, which is different from Fe-free silicates, where small particles are too transparent (cf. Fig.~1 in H{\\\"o}fner 2008). For carbon grains with radii of about $0.1 - 1 \\mu$m, however, the SPL of the Mie theory may severely underestimate the radiative pressure, with possible consequences for the wind properties. The extensive model grid presented by Mattsson et al. (2010, hereafter Paper I) {shows that carbon grains in this size range may be quite common,} in particular for conditions that allow for efficient grain growth (high C abundance, low effective temperatures, slow winds), {indicating} a potential inconsistency with the underlying assumption of SPL opacities.  The main objective of this paper is to establish when the small particle approximation can be applied, and to quantify the possible errors that may occur in mass-loss rates, wind velocities, or dust-to-gas ratios in cases where the dust particles grow beyond this regime. For that purpose, we have implemented a generalized description of dust opacities in our models, using actual mean grain sizes and corresponding radiative cross sections that are valid beyond the SPL (see Sect.~\\ref{s_method}). From the model grid in Paper~I (based on the SPL) we select a subgroup of models that show large grains and/or slow winds, and which we therefore can expect to be noticeably affected by the assumptions about dust opacities. Re-computing these models with the newly implemented, generalized treatment of grain opacities gives estimates of the errors introduced by using the small particle limit. It should be noted, however, that the extreme cases discussed here are not necessarily representative of the majority of the wind models in Paper~I, but that they rather highlight grain size as a potentially critical property, and presumably give an upper limit of the errors. ",
        "conclusions": "In the first paper of this series we presented a large grid of frequency-dependent dynamic models for atmospheres and winds of C-type AGB stars, with the main purpose of providing a more realistic description of dust-driven mass loss as input for stellar evolution models. One of the underlying assumptions of the models, i.e., that dust opacities can be described with the small-particle limit (SPL) of the Mie theory, however, turned out to be questionable with typical emerging grain sizes in a range where the SPL may severely underestimate the actual grain opacities.  Introducing a generalized description of radiative cross sections valid for arbitrary particle radii, we have explored grain size effects on the wind properties of carbon stars. The time-dependent description of grain growth used in our models readily gives particle radii corresponding to various means of the grain size distribution in every layer at every instant of time. To keep the computational effort at a level that will allow the construction of large model grids, we have used descriptions of the dust opacities based on these local mean grain radii.  From the large number of models presented in Paper~I we selected two samples, i.e., a group of models with strong, well-developed outflows and another group close to thresholds for dust-driven winds in stellar parameters space, referred to as \"critical cases\". For the first group, which is presumably representative of most wind-forming models in Paper~I, the effects of grain size on mass-loss rates and wind velocities are small, whereas the critical models show more intense mass loss and (in some cases significantly) higher wind velocities when using the new, generalized description of dust opacities. Models in both groups tend towards lower dust-to-gas ratios, illustrating a self-regulating feedback between grain growth and the increased opacity (and, consequently, higher radiative acceleration) per dust mass. Therefore, in general, the \"dust-loss rates\" can be expected to be lower than in Paper~I.  Extrapolating from the samples investigated here, it seems that the mass loss rates given in Paper~I are reliable, within the limits of current theoretical and observational uncertainties, except for models close to thresholds for dust-driven outflows in stellar parameter space where mass loss is probably underestimated.  Given the results presented here, a full implementation of grain-size-dependent opacities seems to be important as future work. Furthermore, the actual sizes of dust grains may be of great importance for theoretical spectra of dynamic atmosphere models of carbon stars, which we plan to study in the near future."
    },
    "1107/1107.1185.txt": {
        "abstract": "Spitzer/IRS spectra from 5 to 37 $\\mu$m for a complete sample of 31 R Coronae Borealis stars (RCBs) are presented. These spectra are combined with optical and near-infrared photometry of each RCB at maximum light to compile a spectral energy distribution (SED). The SEDs are fitted with blackbody flux distributions and estimates made of the ratio of the infrared flux from circumstellar dust to the flux emitted by the star. Comparisons for 29 of the 31 stars are made with the IRAS fluxes from three decades earlier: Spitzer and IRAS fluxes at 12 $\\mu$m and 25 $\\mu$m are essentially  equal for all but a minority of the sample.  For this minority, the IRAS to Spitzer flux ratio exceeds a factor of three. The outliers are suggested to be  stars where formation of a dust cloud or dust puff is a rare event. A single puff  ejected prior to the IRAS observations may have been reobserved by Spitzer  as a cooler puff at a greater distance from the RCB. RCBs which experience more frequent optical declines have, in general, a circumstellar environment containing puffs subtending a larger solid angle at the star and a quasi-constant infrared flux. Yet, the estimated subtended solid angles and the blackbody temperatures of the dust show a systematic evolution to lower solid angles and cooler temperatures in the interval between IRAS and Spitzer. Dust emission by these RCBs and those in the LMC is similar in terms of total 24 $\\mu$m luminosity and [8.0]$-$[24.0] color index. ",
        "introduction": "The R Coronae Borealis (here, RCB) stars are notable for two distinct peculiarities (Clayton 1996). First, they are hydrogen-poor, helium-rich supergiants: the H-deficiencies range from about 10-100 to at least 10$^8$. Second, the RCB stars experience unpredictable and rapid declines in brightness: declines of 2 to 8 magnitudes in the visual occuring at intervals of less than a year to greater than 20 years last from weeks to months to years. These declines are caused by formation of a cloud of carbon soot above the Earth-facing surface of the star. Discovery of an infrared excess confirmed the obvious suspicion that the stars were dust producers (Stein et al. 1969; Feast et al. 1997). Typical blackbody temperatures of the dust run from about 400 K to 900 K. In a representative case, about one-third of the photospheric flux is absorbed by dust and reemitted in the infrared. The observation that the infrared flux may be little affected by a decline  shows that the dust is distributed in clouds around the star (Forrest et al. 1972).  Recently, high angular resolution images of RY Sgr, a bright RCB, at 2.2, 4.5, and 8$-$13 $\\mu$m showed clearly that the dust is indeed distributed in clouds (de Laverny \\& M\\'{e}karnia 2004; Le\\~ao et al. 2007; Bright et al. 2011). IRAS photometry  at long wavelegnths showed that, in addition to the warm dust, some RCBs have dust at a lower temperature (say, 30 K to 100 K) and, therefore, at large distances from the star (Rao \\& Nandy 1986; Walker 1986; Gillett et al. 1986).  Each of the two principal peculiarities prompts  leading questions. In the case of the H-deficiency, that question is - what are the evolutionary origins of the RCBs that result in a H-poor stars?  Two scenarios remain under active consideration for the RCBs and their putative relatives the H-deficient carbon (HdC) stars to lower temperatures and the extreme helium (EHe) stars to higher temperature. In one, the H-poor supergiant is formed from the merger of a He white dwarf with a C-O white dwarf; the double-degenerate (DD) scenario. In the competing picture, the H-poor supergiant results from a final post-AGB shell flash in the central star of a planetary nebula; the so-called final flash (FF) scenario. In both cases, the trigger - the merger or the final flash - transforms a white dwarf into a H-poor supergiant  for a period of a few thousand years. There is evidence that both the DD and FF scenarios occur but the DD scenario seems  likely to account for the majority of the RCBs.  Several insights into the origins of the RCBs are coming from spectroscopic determinations of the stellar chemical compositions (Lambert \\& Rao 1994; Asplund et al. 2000; Clayton et al. 2005, 2007; Garc\\'{\\i}a-Hern\\'{a}ndez et al. 2009, 2010; Jeffery et al. 2011; Pandey \\& Lambert 2011). Detailed abundance analyses, which are possible for the warm RCBs but rarely undertaken for the cool RCBs with their spectra rich in molecular lines, suggest that many RCBs are  likely fruits of the DD scenario. A few RCBs show several highly unusual abundance signatures and, in particular, very distinctive Si/Fe and S/Fe ratios. Such stars are called `minority' RCBs -  see Lambert \\& Rao (1994) who introduced the terms `majority RCB' and `minority RCB' star.  A rare class of hot RCBs is discussed by De Marco et al. (2002) and includes DY Cen, a minority RCB.  The other principal peculiarity -- the unpredictable declines -- stimulates a series of  questions about the dust around RCB stars such as: What is the composition of the dust? Where does dust form relative to the stellar surface? What triggers  formation of the obscuring cloud?  How frequently do dust clouds form? Does dust form at preferred latitudes on the star or are formation sites spread uniformly across the stellar surface?  Clayton (1996) reviews evidence pertinent to these  questions. Perhaps, the key novel theoretical idea of recent times comes from Woitke et al. (1996) who developed a model in which a pulsation-induced shock triggers dust nucleation near the star (one to two stellar radii out) as  gas behind the outward propagating shock cools below the condensation temperature (say, $< 1500$ K) in the star's upper atmosphere. Presence of cool gas during light minima has now been detected in three RCBs: R CrB (Rao et al. 2006; Rao \\& Lambert 2010), V854 Cen (Rao \\& Lambert 2000), and V CrA (Rao \\& Lambert 2008).  Light variations at maximum light are common among RCBs and generally interpreted as arising from pulsations. Absorption line splitting suggestive of an atmospheric shock is regularly observed for RY Sgr (Danziger 1965; Cottrell \\& Lambert 1982) and occasionally for R CrB. Pugach (1977) noted a correlation between the onset of a light decline for RY Sgr  and the pulsation phase. Crause et al.  (2007) from long-term photometric studies of four RCBs have shown that a decline occurs at a particular phase of the pulsation cycle, although not every pulsation cycle results in a decline. Thus, evidence is accumulating that dust formation occurs near the star.  Radiation pressure on the dust grains is considered to drive them them rapidly outward.  Clues to several of the questions concerning the dust are contained in the shape of the infrared continuum emitted by the circumstellar dust, the presence of emission or absorption features imposed on that continuum, and on the temporal variation of  the infrared flux.  Ground-based infrared spectrophotometry has revealed a smooth continuum in the atmospheric windows; strong  emission and absorption features have not been seen. Inability to observe in regions blocked by the Earth's atmosphere is an especially serious problem in searching for features attributable to dust.  Infrared Space Observatory (ISO) spectra obtained for just three (the brightest) RCBs - R CrB, RY Sgr, and V854 Cen - at a resolution of R=1000 showed for the first time some excess emission over a quasi-blackbody continuum (Lambert et al. 2001). Broad  unidentified emission features were seen centered on about 6 $\\mu$m and 13 $\\mu$m. Emission features from 4 $\\mu$m to 15 $\\mu$m for V854 Cen but not for R CrB and RY Sgr showed a resemblance to a laboratory spectrum of hydrogenated amorphous carbon (Colangeli et al. 1995; Scott et al. 1997).  With the advent of the Spitzer Space telescope, it became possible for the first time to extend  low-resolution infrared spectroscopy to a much larger sample of RCBs. In this paper, we present a library of infrared Spitzer/IRS spectra for a large sample (31) of RCB stars. Spectra are characterized by a fit of blackbodies to optical and  infrared photometry and the Spitzer spectra. Estimates are provided of the infrared flux emitted by dust to the total stellar flux. Comparisons are made with previously reported measurements of the infrared flux, principally the 12$\\mu$m and 25$\\mu$m flux measurements from the IRAS satellite. Infrared emission features superimposed on the blackbody continuua will be discussed in detail in a subsequent paper. Emission spectra of DY Cen and V854 Cen showing features from polycyclic aromatic hydrocarbons (PAHs) and C$_{60}$ have been discussed by Garc\\'{\\i}a-Hern\\'{a}ndez et al. (2011). Spitzer/IRS spectra of the hot RCB stars V348 Sgr and HV 2671 are presented in Clayton et al. (2011).  Section 2 describes our sample of RCB stars, the Spitzer/IRS and  some ground-based photometric observations obtained at about  the same time as the Spitzer observations, and the IRAS observations.  Section 3 constructs the  spectral energy distributions (SEDs) from $\\sim$0.4 to 38 $\\mu$m  from optical and infrared observations, where the SEDs are fitted using blackbodies for the star and the dust for each object. Spitzer and IRAS observations are compared and discussed in Section 4, while a comparison of the RCB dust emission in different metallicity environments is offered in Section 5. Finally, the paper concludes with Section 6. ",
        "conclusions": "Our almost-complete sample of Spitzer/IRS spectra of R Coronae Borealis stars has been combined with multi-color photometry of a RCB at maximum light to provide a spectral energy distribution from, in general, the U or V band to 37 $\\mu$m. Each SED has been fit with a blackbody to represent the stellar flux and one, two, or three blackbodies to represent the emission from the circumstellar dust. A typical RCB emits about 30\\% of the stellar flux in the infrared. Although not discussed in this paper, emission features superimposed on the combination of dust blackbodies are with but a couple of exceptions limited to emission from 6-10 $\\mu$m. The exceptions discussed by Garc\\'{\\i}a-Hern\\'{a}ndez et al. (2011) are DY Cen and V854 Cen where emission is from PAHs and  C$_{60}$ molecules.  For the majority of the RCBs, there is fair agreement between the Spitzer spectrum and the 12 $\\mu$m and 25 $\\mu$m fluxes from IRAS from about three decades previously. There are five exceptions where the IRAS fluxes are between three to five times those recorded by Spitzer. Oddly, these outliers do not have counterparts for which the IRAS fluxes are lower than the Spitzer values.  Our results are consistent with the proposal that clouds of carbon soot form in puffs above the surface of a RCB. There is evidence that puffs are formed randomly and without a prefered direction. For some stars and especially for the above five outliers a single puff may dominate the infrared emission. Our sample of Galactic RCBs and those in the LMC share the same properties of their dusty circumstellar shells, as evidenced by their total 24 $\\mu$m luminosity and [8.0]$-$[24.0] color index.  Although an important next step is to attempt modeling the radiative transfer in the dusty circumstellar environment for our sample of RCB stars, the value of long-term monitoring through infrared photometry and spectroscopy should not be underestimated."
    },
    "1107/1107.4538_arXiv.txt": {
        "abstract": "{Using a suite of simulations \\cite{Gov10} which sucessfully produce bulgeless (dwarf) disk galaxies, we provide an analysis of their associated cold interstellar media (ISM) and stellar chemical abundance patterns. A preliminary comparison with observations is undertaken, in order to assess whether the properties of the cold gas and chemistry of the stellar components are recovered successfully. To this end, we have extracted the radial and vertical gas density profiles, neutral hydrogen velocity dispersion, and the power spectrum of structure within the ISM. We complement this analysis of the cold gas with a brief examination of the simulations' metallicity distribution functions and the distribution of $\\alpha$-elements-to-iron.} ",
        "introduction": "\\label{intro}  \\vspace{-5mm} Historically, simulating the formation and evolution of a bulgeless, rotating, stellar disk, within the context of the classical picture of hierarchical assembly, has proven problematic \\cite{Gov07}.  Overcoming this has only been achieved recently \\cite{Gov10}, with Governato et~al. presenting simulated dwarfs with bulge-to-disk ratios $<$0.04, with very shallow associated central dark matter profiles.  This was achieved by resolving the inhomogeneous ISM, thereby allowing for the imposition of a realistic star formation density threshold which resulted in strong supernovoae-driven outflows preferentially removing low angular momentum gas.  In what follows, we concentrate our analysis on three simulated dwarf galaxies, the basic properties for two of which (DG1, DG1LT) have been introduced previously \\cite{Gov10}. In \\S\\ref{sec:2}, we quantify the spatio-kinetic characteristics of these bulgeless dwarfs, and contrast them with recent observational work \\cite{Obr10,Stan99,Tam09}.  A preliminary analysis of the chemical properties of the associated stellar components is then presented in \\S\\ref{sec:3} and compared with relevant empirical work in the field \\cite{Scho09,Tol09}. As the simulations are not modelled on any one particular galaxy, the analysis should be seen as comparative only in a qualitative sense.   \\subsection{Simulations} \\label{simulations}  \\vspace{-5mm} The three simulations analysed were run using the N-body$+$SPH code \\textsc{gasoline} \\cite{Wad04}, employing the same initial conditions, and differing only in the adopted resolution, star formation density threshold, feedback efficiency, and star formation efficiency. The base simulation (DG1) \\cite{Gov10} uses a star formation density threshold a thousand times higher (100~cm$^{-3}$) than that adopted in previous cosmological simulations ($\\sim$0.1~cm$^{-3}$); a lower density threshold version (DG1LT), despite its possession of classical problematic traits (i.e., centrally-concentrated light (B/D$\\approx$0.3) and dark matter density distributions), is included here for comparison.  This implementation of a more realistic density threshold for star formation was made possible by the associated high force resolution (86pc) which allowed us to resolve individual star forming regions of mass $\\sim$10$^5$~M$_\\odot$.  We include in our analysis a new simulation (nDG1), similar to DG1, but with two important modifications - the inclusion of high-temperature (T$>$10$^4$K) metal-line cooling \\cite{SWS10}, and enhanced supernova feedback (100\\% of the energy being coupled to the surrounding ISM in the form of thermal energy, rather than the 40\\% used in the base simulation). The feedback method is described fully in \\cite{Stinson06}. ",
        "conclusions": ""
    },
    "1107/1107.3404_arXiv.txt": {
        "abstract": "{ARGO-YBJ is a \\emph{full coverage} air shower array located at the YangBaJing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$) recording data with a duty cycle $\\geq$85\\% and an energy threshold of a few hundred GeV. In this paper the latest results in Gamma-Ray Astronomy are summarized. ",
        "introduction": " ",
        "conclusions": "The ARGO-YBJ experiment is an air shower array with large field of view and has been continuously monitoring the northern sky since November 2007. With all data recorded up to February 2011, 4 sources have been observed with a statistical significance greater than 5 s.d.: Crab Nebula, Mrk421, MGRO J1908+06 and MGRO J2031+41.  ARGO-YBJ has been continuously monitoring the Crab Nebula for more than 3 years. The measured energy spectrum in the energy range 500 GeV - 10 TeV is in agreement with the results of other experiments. No flux variations with a statistical significance larger than 5 s.d. have been detected in time scales of days or months. Nevertheless, enhanced flux of significance about 3 s.d. has been observed in coincidence with the occurrence of two flares detected by AGILE and Fermi.  In the last 3 years we carried out a detailed long-term monitoring of Mrk421 detecting a number of TeV $\\gamma$-ray flares correlated with X-ray flaring activity.  The energy spectra of MGRO J1908+06 and MGRO J2031+41 have been measured. No signal from MGRO J2019+37 is detected, and the derived upper limits at 90\\% c.l. are lower than the Milagro flux at energies below 3 TeV, indicating that the source flux may be not stable."
    },
    "1107/1107.5825_arXiv.txt": {
        "abstract": "\\noindent Ices regulate much of the chemistry during star formation and account for up to 80\\% of the available oxygen and carbon. In this paper, we use the {\\it Spitzer} $c2d$ ice survey, complimented with data sets on ices in cloud cores and high-mass protostars, to determine standard ice abundances and to present a coherent picture of the evolution of ices during low- and high-mass star formation. The median ice composition H$_2$O:CO:CO$_2$:CH$_3$OH:NH$_3$:CH$_4$:XCN is 100:29:29:3:5:5:0.3 and 100:13:13:4:5:2:0.6 toward low- and high-mass protostars, respectively, and 100:31:38:4:--:--:-- in cloud cores.  In the low-mass sample, the ice  abundances with respect to H$_2$O of CH$_4$, NH$_3$, and the component of CO$_2$ mixed with H$_2$O  typically vary by $<$25\\%, indicative of co-formation with H$_2$O. In contrast, some CO and CO$_2$ ice components, XCN and CH$_3$OH vary by factors 2--10 between the lower and upper quartile.  The XCN band correlates with CO, consistent with its OCN$^{-}$ identification. The origin(s) of the different levels of ice abundance variations are constrained by comparing ice inventories toward different types of protostars and background stars, through ice mapping, analysis of cloud-to-cloud variations, and ice (anti-)correlations. Based on the analysis, the first ice formation phase is driven by hydrogenation of atoms, which results in a H$_2$O-dominated ice. At later prestellar times, CO freezes out and variations in CO freeze-out levels and the subsequent CO-based chemistry can explain most of the observed ice abundance variations. The last important ice evolution stage is thermal and UV processing around protostars, resulting in CO desorption, ice segregation and formation of complex organic molecules. The distribution of cometary ice abundances are consistent with with the idea that most cometary ices have a protostellar origin. ",
        "introduction": "\\noindent Star formation begins with the collapse of an interstellar cloud core to form a protostar, which grows from infall of envelope material and later through an accretion disk, where planets may  form. This physical evolution from cores to planetary systems is accompanied by a chemical evolution, which will affect planet and planetesimal compositions. Much of this chemical evolution takes place in icy mantles on interstellar grain surfaces, and the aim of this paper is to consolidate the knowledge of ice evolution during star formation coming out of the {\\it Spitzer} Space Telescope mission.  The first ices were detected in the interstellar medium almost 40 years ago \\citep{Gillet73} and from previous work, mostly toward high-mass protostars, H$_2$O, CO and CO$_2$ ices are known to be common during the cold and dense stages of star formation, with abundances reaching 10$^{-4}$ $n_{\\rm H_2}$. The ice mantles also contain smaller amounts of CH$_3$OH, CH$_4$, NH$_3$ and XCN \\citep[e.g][]{Gibb04}. CH$_3$OH and other ices are proposed sources of complex organic molecules \\citep{Charnley92,Garrod08,Oberg09d} and determining ice abundances and production channels in star forming regions is therefore of great interest for studies of prebiotic chemistry. Most ices except for CO are predicted to form {\\it in situ} on interstellar grain surfaces through hydrogenation and oxygenation of atoms and small molecules \\citep{Tielens82}. Heat and UV radiation from protostars may result in additional ice processing and \\citet{Gibb04} suggested that some CO$_2$ ice, and all CH$_3$OH and XCN ice form through such processing based on observations toward high-mass protostars with the {\\it Infrared Space Observatory} (ISO). This has been challenged by observations of abundant CH$_3$OH and XCN ice toward low-mass protostars \\citep{Pontoppidan03,vanBroekhuizen05}, where ices are protected from strong UV fields during most of their life time.  \\textit{Spitzer}'s high sensitivity made it possible to observe ices toward low mass protostars and background stars at 5--30~$\\mu$m. The $c2d$ program \\citep{Evans03} obtained IRS spectra of 50 low-mass protostars \\citep[][from now on Paper I-IV]{Boogert08, Pontoppidan08, Oberg08, Bottinelli10}, providing an unprecedented sample size that complements the {\\it ISO} data on high-mass protostars. Additional {\\it Spitzer} ice observations exist on ices toward background sources, looking through molecular clouds at a range of extinctions \\citep{Bergin05,Knez05,Pontoppidan06,Whittet09,Boogert11} and toward protostars in a high-UV environment \\citep{Reach09}.  Because the  {\\it Spitzer} spectrometer had a lower cutoff at $\\lambda \\sim 5$~$\\mu$m, complete ice inventories can only be obtained by adding complementary ground-based spectroscopy to cover the strong 3 $\\mu$m H$_2$O, the 4.65 $\\mu$m CO, the 4.6 $\\mu$m XCN, and the 3.53 $\\mu$m CH$_3$OH features. The latter feature has been used to validate the use of the 9.7 $\\mu$m feature to derive CH$_3$OH ice abundances (Paper I). The XCN feature has been shown empirically to consist of two components, whose relative contributions to the feature vary from source to source \\citep{vanBroekhuizen05}. One of the components can be securely assigned to OCN$^-$ from comparison with laboratory spectra, while the carrier of the second component, peaking at 2175 cm$^{-1}$, is unknown. Suggested carriers of the 2175 cm$^{-1}$ component include OCN$^-$ present in a different ice environment, resulting in a different spectrum, as well as completely different molecules \\citep{Pendleton99,vanBroekhuizen04}.  In the analysis of astrophysical ice spectra, comparison with laboratory ice spectra is key, both to identify ice species (such as OCN$^-$ above) and to characterize the ice morphology; i.e., some ice bands, such as CO and CO$_2$, cannot be fitted by a single laboratory ice mixture, but seem to trace molecules in two or more ice phases \\citep[][Paper II]{Tielens91, Chiar95,Dartois99}. Traditionally the observed spectra have been directly compared to a superposition of pure and mixed ice spectra, characterizing the ice morphology in each source independently \\citep[e.g.][]{Merrill76,Gibb04, Zasowski09}. The constraints are often degenerate, however, since ice spectral features vary with ice composition, temperature and radiation processing. To address this, ice bands have instead been decomposed phenomenologically into a small set of unique spectral components \\citep[e.g.][ and \\S2.2]{Tielens91, Pendleton99, Keane01,Pontoppidan03b}. The contributions of the derived components are then used to characterize the spectral bands in all sources in the sample. This method was employed in papers I and II. Because all observed spectra are decomposed into the same, small number of components this method provides information on the sample as a whole, i.e. it directly shows which parts of the spectral profile are ubiquitous and which are environment dependent. This is crucial information when assigning a component carrier -- without this, the degeneracy is almost always too large to draw conclusions about the structure of the ice from spectral profiles alone.  Building on the analysis in Paper I-IV, this paper aims to establish an ice evolution scenario that also takes into account the present knowledge of ice abundances toward cloud cores \\citep{Boogert11} and high-mass protostars \\citep{Gibb04}. Section \\ref{sec:obs} summarizes the sample characteristics. Section \\ref{sec:res} presents new ice data for 10 low-mass protostars, followed by ice abundance medians toward low- and high mass protostars. The variation among the low-mass protostellar abundances are investigated followed by comparison of this ice sample with ices toward high-mass protostars, low-mass protostars in a high-radiation environment, and background stars. The reasons for the abundance variations seen for some ice features are further explored through analysis of spatial differences and correlation studies. Further, new data are used to test previous hypotheses about the carrier(s) of the XCN band. The results are discussed in \\S \\ref{sec:disc} with respect to different ice formation scenarios and ice chemistry in low-mass versus high-mass star-forming regions. ",
        "conclusions": "\\noindent Large samples of ice sources spanning diverse environments, evolutionary stages and luminosities are necessary both to determine `typical' ice abundances and to identify how ice processes depend on their environment. The main findings from combining the $c2d$ and other ice surveys carried out in a homogenous way are listed below:  \\begin{enumerate} \\item  The CO and CO$_2$ abundance medians relative to H$_2$O are both $\\sim$29\\% toward low-mass YSOs, while CH$_4$, NH$_3$ and CH$_3$OH are more than a factor of five less abundant. \\item CO$_2$:H$_2$O, CH$_4$ and NH$_3$ abundances vary little with respect to H$_2$O, suggesting co-formation. In contrast, CO:H$_2$O, OCN$^-$, CO$_2$:CO and CH$_3$OH vary by factors 2--3 (lower to upper quartile) with respect to H$_2$O, indicative of a separate formation pathway from H$_2$O ice. Pure CO and CO$_2$ ice abundances vary even more, consistent with their sensitivity to protostellar heating. \\item Ice abundances toward background stars and protostars are similar except for a lack of features associated with ice heating, such as pure CO$_2$ ice, toward background stars. In particular, there is no evidence for a different range of CH$_3$OH and CO$_2$ ice abundances. \\item Protostellar ice abundances in nearby star-forming regions do not vary significantly between different clouds, except for a significantly lower pure CO abundances toward the protostars in Taurus. \\item Compared to low-mass YSOs, ice abundances toward high-mass YSOs are characterized by low levels of CO and CO$_2$ ice, which can be explained by ice heating. Low-mass protostars in the highly irradiated region IC 1396A have typical low-mass protostellar CO$_2$ ice abundances, confirming that the original ice conditions during low- and high-mass star formation are similar. \\item Correlation studies within the low-mass protostellar sample show a close association between CO, CO$_2$:CO, CO:H$_2$O and the XCN band, supporting the latter's identification with OCN$^-$. \\item Combining the above information, ice formation can generally be divided into three stages: an early phase driven by atomic hydrogenation reactions in clouds, which are fast compared to cloud core formation time scales of $\\sim$10$^5$ years;  a later CO-freeze-out dependent ice formation phase which takes place in the pre-stellar phase; and finally the protostellar phase where thermal and possibly UV processing shape the ice content. \\item Toward both low- and high-mass protostars and toward molecular clouds, an average of 8--16\\% of the total C, O and N are bound up in ices. Toward the most ice-rich source (IRS 51) 60--80\\% of the non-refractory C and O are in ices. \\item There is evidence for more complex ices toward both low- and high-mass protostars, e.g. HCOOH, CH$_3$CHO and/or C$_2$H$_5$OH, but high-resolution spectra toward more sources are required to confirm their presence and quantify their abundances. \\end{enumerate}  {\\it Facilities:} \\facility{Spitzer, VLT, Keck}"
    },
    "1107/1107.2258_arXiv.txt": {
        "abstract": "s{Model-independent parametrisations of modified gravity have attracted a lot of attention over the past few years; numerous combinations of experiments and observables have been suggested to constrain these parameterisations, and future surveys look very promising. Galaxy Clusters have been mentioned, but not looked at as extensively in the literature as some other probes. Here we look at adding Galaxy Clusters into the mix of observables and examine whether they could improve the constraints on the modified gravity parameters. In particular, we forecast the constraints from combining the Planck CMB spectrum and SZ cluster catalogue and a DES-like Weak Lensing survey. We've found that adding cluster counts improves the constraints obtained from combining CMB and WL data. } ",
        "introduction": "Einstein's General Relativity (GR) is one of the principal ingredients of modern cosmology. Nonetheless, it is our job as physicists to continue to test even the fundamental pillars of cosmology in order to refine, improve and further justify our model of the universe. Testing GR outside of the solar system can be quite challenging, particularly as the effects of a different theory of gravity could be degenerate with different possible constituents of the universe. This is the case with the current cosmological observations that suggest the presence of dark matter and dark energy. As well as explaining these observations, there are also fundamental physics reasons for considering different theories of gravity: GR is inconsistent with quantum mechanics and the search for \"Quantum Gravity\" is one of the holy grails of modern physics. \\\\ Here, we are interested in testing deviations from GR in a model independent way. There are several advantages to a model independent approach; some alternatives to GR do exist but there is no complete theory of, for example, quantum gravity to draw on. Also, amongst the many options there are no \"stand-out\" candidates that are universally considered to be strong alternatives. If a model independent approach suggests that the data is inconsistent with concordance cosmology and GR, it will be relatively unambiguous and therefore a strong motivator to develop alternative theories, as well as possibly giving us a clue as to the nature of these theories.\\\\ There are studies in the literature on the constraining power of current data \\cite{gbz10,bean10} and the general conclusion is that the concordance cosomology is consistent with all of the current data, although the data isn't strongly constraining. Work has also gone into forecasting future experiments \\cite{serra09,gbz09} and again there is a fair degree of consensus here, namely that future surveys will greatly improve prospects. In this work we will examine the constraints that can be put on model independent modified gravity using the combination of Cosmic Microwave Background anisotropies (CMB) cross-correlated with weak lensing surveys and galaxy cluster counts. For more details and up to date figures, see \\cite{me}. ",
        "conclusions": "Over the next 5-10 years, deviations from GR should be well constrained, and the concordance cosmology will either be more secure or may even have undergone a paradigm shift. If the latter is the case, then the results from the model independent tests could be crucial in helping to find a new theory of gravity."
    },
    "1107/1107.0294_arXiv.txt": {
        "abstract": "Short-lived radioisotopes (SLRI) such as $^{60}$Fe and $^{26}$Al were likely injected into the solar nebula in a spatially and temporally heterogeneous manner. Marginally gravitationally unstable (MGU) disks, of the type required to form gas giant planets, are capable of rapid homogenization of isotopic heterogeneity as well as of rapid radial transport of dust grains and gases throughout a protoplanetary disk. Two different types of new models of a MGU disk in orbit around a solar-mass protostar are presented. The first set has variations in the number of terms in the spherical harmonic solution for the gravitational potential, effectively studying the effect of varying the spatial resolution of the gravitational torques responsible for MGU disk evolution. The second set explores the effects of varying the initial minimum value of the Toomre $Q$ stability parameter, from values of 1.4 to 2.5, i.e., toward increasingly less unstable disks. The new models show that the basic results are largely independent of both sets of variations. MGU disk models robustly result in rapid mixing of initially highly heterogeneous distributions of SLRIs to levels of $\\sim$ 10\\% in both the inner ($<$ 5 AU) and outer ($>$ 10 AU) disk regions, and to even lower levels ($\\sim$ 2\\%) in intermediate regions, where gravitational torques are most effective at mixing. These gradients should have cosmochemical implications for the distribution of SLRIs and stable oxygen isotopes contained in planetesimals (e.g., comets) formed in the giant planet region ($\\sim$ 5 to $\\sim$ 10 AU) compared to those formed elsewhere. ",
        "introduction": "The short-lived radioisotope (SLRI) $^{60}$Fe appears to have been synthesized in a Type II supernova (Mostefaoui, Lugmair, \\& Hoppe 2005; Tachibana et al. 2006) and injected into the presolar cloud (Boss et al. 2008, 2010; Boss \\& Keiser 2010) from the same massive star that is likely to be the source of the bulk of the solar nebula's $^{26}$Al (Limongi \\& Chieffi 2006; Sahijpal \\& Soni 2006). Given the injection of SLRIs into the presolar cloud by Rayleigh-Taylor fingers (Boss et al. 2008, 2010; Boss \\& Keiser 2010), it might be expected that the SLRIs would be initially highly spatially and temporally heterogeneous in their distribution in the solar nebula. However, the nearly identical Fe and Ni isotopic compositions of iron meteorites, chondrites, and the Earth require that the injected $^{60}$Fe must have been mixed to less than 10\\% heterogeneity in the solar nebula (Dauphas et al. 2008). A similar constraint arises from the need to preserve the use of $^{26}$Al as an accurate nebular chronometer (e.g., Thrane, Bizzarro, \\& Baker 2006), while simultaneously allowing for the spread of stable oxygen isotope ratios (Lyons \\& Young 2005; Lee, Bergin, \\& Lyons 2008). Three-dimensional hydrodynamical models of the evolution of a marginally gravitationally unstable (MGU) solar nebula have shown that mixing of initially highly heterogeneous distributions of SLRIs can indeed reduce the level of heterogeneity to $\\sim$ 10\\% or lower in less than 1000 yrs (Boss 2004a, 2006, 2007, 2008).  The discovery of refractory grains among the particles collected from Comet 81P/Wild 2 by the Stardust spacecraft (Brownlee et al. 2006; Simon et al. 2008; Nakamura et al. 2008) provided the first ground truth for large-scale transport of materials formed in high temperature regions close to the protosun outward to the comet-forming regions of the solar nebula. One refractory particle found by Stardust, Coki, has an age $\\sim$ 1.7 Myr younger (Matzel et al. 2010) than that of calcium, aluminum-rich inclusions (CAIs), implying that outward radial transport continued for millions of years after CAI formation. Measurements of the iron sulfide content of particles from Wild 2 imply that over half of the comet's mass derived from the inner solar nebula (Westphal et al. 2009). Similar hydrogen, nitrogen, and oxygen isotopic anomalies occur in both primitive meteorites and in cometary dust particles, implying that meteorites and comets both formed from the same basic mixture of disk material (Busemann et al. 2009; Al\\'eon et al. 2009).  Observations of disks around young stars often find evidence for crystalline silicate grains at distances ranging from inside 3 AU to beyond 5 AU, in both the disk's midplane and its surface layers (e.g., Mer\\'in et al. 2007). Only the most massive AGB stars produce significant quantities of crystalline grains (Speck et al. 2008); grains in the interstellar medium are primarily amorphous. Crystalline silicate grains could have been produced through thermal annealing of amorphous grains (e.g., Sargent et al. 2009a,b) by the hot disk temperatures reached only within the innermost disk, well inside 1 AU. Crystalline and amorphous silicate grains in most protoplanetary disks may thus have experienced the same large-scale transport phases as the refractory particles found in Wild 2.  While accretion disk models driven by a generic turbulent viscosity have long been invoked as a means to explain large-scale transport (Gail 2001, 2002, 2004; Tscharnuter \\& Gail 2007; Ciesla 2007, 2008, 2009, 2010a,b; Birnstiel et al. 2009; Hughes \\& Armitage 2010; Heinzeller et al. 2011; Jacquet et al. 2011), the detailed physics behind $\\alpha$-viscosity remains unclear, especially considering that the magneto-rotational instability (MRI) often assumed to be the source of the $\\alpha$-viscosity is unable to drive disk evolution in the magnetically dead midplane regions (e.g., Matsumura \\& Pudritz 2006) of most interest for planetary formation. Objections have also been raised to the assumption that angular momentum transport in a MRI-driven disk can be described by the standard model for $\\alpha$-viscosity (Pessah, Chan, \\& Psaltis 2008). In contrast, a MGU disk presents a self-consistent mechanism for studying mixing and transport in protoplanetary disks, with no free parameters beyond the initial choice of a disk massive enough, and cold enough, to be MGU.  We present here a new set of three-dimensional MGU disk models similar to those studied previously (e.g., Boss 2004a, 2006, 2007, 2008), but with several variations intended to test the robustness of the conclusions about mixing and transport of the previous models. Given that a MGU disk appears to be a likely requirement for the formation of gas giant planets, by either core accretion (e.g., Inaba et al. 2003; Chambers 2008) or by disk instability (e.g., Boss 2010), these MGU models are intended to learn what cosmochemical consequences might also derive from such phases of rapid disk evolution. ",
        "conclusions": "All eight of these new models evolve in much the same manner as model 9S of Boss (2008): the initially high degree of SLRI heterogeneity is lowered by mixing within $\\sim$ 500 yrs to a dispersion of $\\sim$ 10\\% inside 5 AU and $\\sim$ 2\\% from 5 AU to 10 AU. Combined with previous results for disks extending to 20 AU (Boss 2007), yielding dispersions of $\\sim$ 10\\% beyond 10 AU, gradients in isotopic heterogeneity are to be expected in MGU disks. MGU disk models thus make a clear prediction: solids that accreted in the gas giant planet region from $\\sim$ 5 to 10 AU should be significantly more isotopically homogenized than those that formed interior to, or exterior to, this region. This bold prediction must be taken with a grain of salt, however, as it presumes locally closed system behavior, and does not take into account other factors, such as the arrival of multiple Rayleigh-Taylor fingers with varying time intervals, or ongoing photochemistry due to self-shielding of CO molecules at the disk's surfaces.   Given that all meteoritical samples are believed to originate from planetestimals formed in the inner solar nebula, there may not be any measurable cosmochemical consequence of this predicted gradient in isotopic heterogeneity. Nevertheless, the robustness of the present MGU disk models supports the idea of the existence of a low degree of spatial heterogeneity of SLRIs in the inner solar nebula, consistent with the long-held belief that cosmochemical variations in inferred initial $^{26}$Al/$^{27}$Al ratios can be used as accurate chronometers for the timing of planetesimal and planet formation events in the inner solar system. While the FUN inclusions show no evidence for live $^{26}$Al and so are thought to have formed prior to the arrival of the first Rayleigh-Taylor finger (Sahijpal \\& Goswami 1998), most refractory inclusions have $^{26}$Al/$^{27}$Al ratios that scatter by no more than $\\sim$ 10\\% around the mean value of $\\sim 4.5 \\times 10^{-5}$ (MacPherson et al. 1995; Young et al. 2005). This level of scatter in $^{26}$Al/$^{27}$Al ratios appears to be consistent with the 10\\% dispersion expected in the inner regions of disks that have experienced a MGU phase.  However, there may well be a cosmochemically observable signature of this predicted isotopic heterogeneity gradient in comets, and in the interplanetary dust particles (IDPs) that originate in comets. Comets that formed in the region of the gas giant planets are believed to have been kicked outward by close encounters with the giant planets, either to hyperbolic orbits, or to highly eccentric orbits, i.e., to the Oort Cloud (e.g., Paulech et al. 2010). Long-period comets are believed to be derived from the Oort Cloud and thus to have originated from the most isotopically homogeneous region of a MGU disk. Short-period comets, on the other hand, are thought to be derived from the Kuiper Belt, i.e., from an origin in the outer disk (e.g., Levison et al. 2008), where the same high degree of isotopic homogeneity would not have occurred in a MGU disk. Assuming that these ideas about the origins of long-period and short-period comets are correct, MGU disk models make a clear prediction: long-period comets should be more isotopically homogeneous than short-period comets. In addition, both types of comet should contain refractory particles similar to CAIs, such as the Coki particle in Wild 2."
    },
    "1107/1107.3154_arXiv.txt": {
        "abstract": "Efforts are being made to observe the 21-cm signal from the `cosmic dawn' using sky-averaged observations with individual radio dipoles. In this paper, we develop a model of the observations accounting for the 21-cm signal, foregrounds, and several major instrumental effects. Given this model, we apply Markov Chain Monte Carlo techniques to demonstrate the ability of these instruments to separate the 21-cm signal from foregrounds and quantify their ability to constrain properties of the first galaxies. For concreteness, we investigate observations between 40 and 120 MHz with the proposed {\\it DARE} mission in lunar orbit, showing its potential for science return. ",
        "introduction": "\\label{sec:intro}  One of the remaining frontiers of modern cosmology is the end of the `dark ages' and the `cosmic dawn'.  This is the period ranging from roughly 100 million years ($z\\sim30$) to a billion years ($z\\sim6$) after the Big Bang, when the first stars and galaxies formed, lighting up the Universe.  This period lies at the edge of current observational techniques and is of considerable theoretical interest. The next decade is expected to see significant improvements in observations as telescopes such as the {\\em James Webb Space Telescope} ({\\it{}JWST}) and the Atacama Large Millimeter Array (ALMA) go online.  These instruments will provide considerable information about galaxy formation at $z\\lesssim 10$--$15$, but even these large telescopes will be hard pressed to probe the very beginning of the cosmic dawn.  Measurements of 21-cm emission and absorption from intergalactic hydrogen at high redshift promise to increase greatly our knowledge of the Universe at redshifts $z\\gtrsim 6$ \\citep*{madau1997}. Experiments under way at present are concentrating on frequencies $\\nu\\gtrsim 100\\ \\mathrm{MHz}$ ($z\\lesssim 13.2$), and are hoping to capture the transition from an almost completely neutral Universe to an almost completely ionized one (the epoch of reionization, or EoR). Future observations at yet lower frequencies (higher redshifts) may probe the epoch when the first sources formed -- `cosmic dawn' -- and even the preceding `dark ages'.  There are two main approaches to making these measurements: using a large interferometric array to produce statistics (e.g.\\ power spectra), and perhaps even images, of the 21-cm brightness temperature; or using a single antenna to measure the mean brightness temperature as a function of frequency and redshift \\citep{SHA99}. In either case, the bright foregrounds at low frequencies present one of the most significant challenges to extracting the 21-cm signal. The difficulty is alleviated somewhat in the former approach since an interferometer is sensitive only to fluctuations in the foregrounds, which are small compared to the mean on the scales of interest, but they still exceed the 21-cm fluctuations in intensity by several orders of magnitude. Interferometric measurements have other benefits too. For example, the spectrum of fluctuations carries more information than the mean signal alone, and interferometers may make it easier to identify and excise man-made radio-frequency interference (RFI).  An interferometer cannot measure the mean signal, however. Moreover, global signal experiments designed to measure the mean brightness temperature may be much simpler and cheaper than large arrays, and are not troubled to the same extent by distortions caused by the Earth's ionosphere.  The large sky temperature at these frequencies also means that the sky makes the dominant contribution to the system temperature, and hence to the sensitivity of the observation for a given bandwidth and integration time. The brightness temperature, $T_\\mathrm{B}$, of the diffuse foregrounds depends on the observing frequency, $\\nu$, as $T_\\mathrm{B}\\sim\\nu^{-2.5}$ \\citep{ROG08}, so interferometric measurements during the cosmic dawn at less than 100 MHz require very long integration times or arrays with a very large collecting area. Because ionospheric effects also become more serious at low frequencies, it has been suggested that the far side of the Moon, which is also free (as yet) from RFI, would be the best and perhaps the only site for an array to probe the cosmic dawn and dark ages \\citep[e.g.][]{burns1988,burns2009,jester2009}.  Building and operating such an array of the requisite size would be quite a formidable undertaking, so global signal experiments provide the best hope for probing the 21-cm signal at $z\\gtrsim 15$ in the near future.  The EDGES experiment \\citep{BOW10}, operating at 100--200 MHz, has pioneered global 21-cm measurements, recently placing limits on how rapidly the global 21-cm signal may vary with frequency, and thereby putting a lower limit on the duration of the reionization epoch.  Even from its superb radio-quiet site in Western Australia, however, it encountered RFI from sources such as telecommunications satellites and radio and television transmitters. The signals from these may reach EDGES quite directly, or arrive via e.g.\\ tropospheric scattering or reflections from aircraft and meteor trails.  This requires a large fraction of the data to be discarded, which would be more damaging at low frequencies where longer integrations are required, and it imposes stringent demands on the dynamic range of the receiver.  A dipole antenna in orbit around the Moon could avoid these problems, since it would be free of RFI when shielded from the Earth over the lunar farside. In addition, an antenna in space experiences a simpler and more stable environment than one on the Earth's surface, which may allow for more straightforward calibration. The use of lunar orbit does not require anything to be landed on the Moon's surface, unlike for a farside array. Such a mission concept has been developed, called the {\\it Dark Ages Radio Explorer} \\citep[{\\it DARE};][]{burns2011}\\footnote{http://lunar.colorado.edu/dare/}. In this paper we therefore explore the constraints that a mission such as {\\it DARE} could place on a model of the 21-cm brightness temperature between the end of the cosmic dark ages and the start of the epoch of reionization. We aim to include all the most important contributions to the low-frequency radio spectra measured by a dipole in lunar orbit: the redshifted 21-cm signal itself; spatially varying diffuse foregrounds based on an empirical model of the low-frequency radio sky; the Sun; the thermal emission of the Moon; the reflection of emission from other sources by the Moon; the response of the instrument, which is based on an electromagnetic model of an antenna design proposed for the {\\it DARE} mission; and the thermal noise for a realistic mission duration. Parameters describing all these components must be fit simultaneously from the data since, for example, it may be that the properties of the instrument cannot be computed or measured on the ground with sufficient accuracy to allow recovery of the 21-cm signal in the presence of the very bright foregrounds.  We therefore extend the work of \\citet{PRI10}, who used Fisher matrix and Monte Carlo methods to predict the accuracy with which models of the 21-cm signal could be constrained by a single antenna, but who considered the simpler case of an experiment which measured a single, deep spectrum (i.e.\\ they did not consider the variation of foregrounds over the sky), and where the only contributions to the measured spectrum were the redshifted 21-cm signal, diffuse foregrounds and noise.  The techniques that we develop, and the basic form of our model for the 21-cm global signal, are quite generic and may be applied to future experiments both on the ground and in space.  For concreteness, we focus here on the proposed {\\it DARE} mission, but a similar methodology could be applied to EDGES and other ground based experiments.  We plan to investigate this in the near future.  We start by outlining the relevant features of our reference experiment, a proposed mission to measure the 21-cm global signal from lunar orbit, in Section~\\ref{sec:mission}. Then, in Section~\\ref{sec:sims} we describe all the different effects which are included in our simulations of data from such a mission, including the parametrizations we use. We also discuss some other contributions, such as impacts of exospheric dust on the antenna and radio recombination lines, and justify neglecting them in this analysis.  Constraints on the model parameters are derived using a Markov Chain Monte Carlo (MCMC) method. In Section~\\ref{sec:fitting} we introduce our implementation of this technique and, as an example, show how well the parameters are recovered by a perfect instrument, which is sensitive across the whole frequency band and whose properties are known exactly. In Section~\\ref{sec:disc} we consider a more realistic instrument with an imperfectly known response, and look at the impact of the various processes we model on the quality of our constraints. Finally, we offer some conclusions in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have presented a model for the data from a proposed lunar-orbiting satellite to measure the global, redshifted 21-cm signal between 40 and 120 MHz. Fitting the parameters of this model to a realistic simulated data set using an MCMC algorithm yields constraints on the 21-cm signal that are comparable to those found using much simpler models for the foregrounds and instrument, despite the fact that we use the data to constrain $\\approx\\! 73$ parameters, rather than 10. The key assumptions used in extracting the signal are that the foregrounds are smooth, that the instrumental response is also smooth and can be determined reasonably well by independent measurements, and that the 21-cm signal, averaged over the solid angle of our antenna beam, is constant across the sky while the foregrounds are not.  A mission of reasonable duration ($\\sim 3$ yr) can find the position of the bottom of the `cosmic dawn' absorption trough in our fiducial model with an accuracy of around $\\pm 1\\ \\mathrm{MHz}$ in frequency and $\\pm 20\\ \\mathrm{mK}$ in temperature (2-$\\sigma$ errors), provided that the instrumental response has been well characterized. The frequency position of the peak in emission at the onset of reionization can be found to within $\\pm 0.5\\ \\mathrm{MHz}$, while `turning point B', marking the onset of Ly$\\alpha$ pumping, can be determined with a 1-$\\sigma$ error of around $\\pm 2.5\\ \\mathrm{MHz}$. For a shorter mission, of e.g.\\ 1000 h, these constraints degrade somewhat, and it may only be possible to find an upper limit on the frequency of turning point B. A mission of 10000 h allows a good measurement of the frequency of turning point B, with a 2-$\\sigma$ confidence interval that lies entirely within the {\\it DARE} frequency band.  We have examined the effect of using prior information on the non-diffuse foregrounds, which may be amenable to measurement from the ground, and on the instrumental parameters. Priors on the foregrounds do not help a great deal, though clearly it will still be valuable to have independent, ground-based measurements of the foregrounds, to inform our modelling and to check that our measurements are consistent.  Tightened priors on the instrument model, however, which correspond to improved calibration, reduce the statistical errors. They may also help to reduce the importance of temperature errors which are correlated across the frequency band, and which result in an uncertainty in the overall normalization of the 21-cm brightness temperature. Even if this correlation is present, it is likely that the shape of the 21-cm signal can be recovered accurately, since the absolute error on the temperature of each of the three turning points tends to be similar.  Interferometric experiments must also perform foreground subtraction, to an accuracy of around one part in $10^3$ for the diffuse Galactic emission, and to one part in $10^6$ or even $10^8$ for bright point sources \\citep*{datta2010}. In this sense, arrays such as MWA, the Low Frequency Array (LOFAR) and the Precision Array to Probe the Epoch of Reionization (PAPER) should characterize the properties of the Galactic and extragalactic foregrounds and validate the assumption that they can be modelled using functions which deviate little from power laws. They differ from experiments such as {\\it DARE} or EDGES, however, in that they will use observations of specific point sources for calibration of gain and bandpass in a way that is not possible for sky-averaged experiments, since the latter cannot isolate the contribution to the measured spectrum from an individual source. For this reason, upcoming arrays have not been designed to achieve an intrinsically smooth bandpass that can be quantified with only a few parameters, and therefore will likely shed little light on the calibration or instrument modelling for sky-averaged experiments.  In this paper, we have focused on a particular reference experiment to illustrate our techniques.  The methodology developed here is very general and can easily be extended to other global 21-cm experiments. As global 21-cm experiments continue to improve from their current relative infancy, there will be a need for improved techniques of statistical analysis. We have taken some early steps in that direction.  It is worth reiterating, however, that there are several other stages in the data analysis which must be passed before the methodology of this paper can be applied. Individual spectra taken with a short cadence (of e.g.\\ 1~s) must be combined together using a map-making procedure to produce something like the eight independent spectra seen here. The frequency response must be internally calibrated, for example by toggling the receiver input between the antenna feeds and calibration loads. Narrow features such as RRLs or, in the case of ground-based experiments, RFI, must be excised. All these steps become more complicated for ground-based experiments. For the map-making, an experiment fixed to the ground would not have the complete control over the pointing direction provided by a satellite, and would not have access to the whole sky. Moreover, the ionosphere effectively causes the sky seen by the antenna to vary with time. Internal calibration is made more awkward by changes in temperature and atmospheric conditions, while a space environment is more predictable. Finally, RFI is likely to be considerably more prevalent than RRLs. The main effects of these earlier steps on the MCMC method are likely to be the introduction of non-Gaussianity to the noise on the frequency spectra, and correlation between different sky areas, both of which affect the computation of the likelihood. It is not clear whether some of these effects could be captured with extra nuisance parameters in the MCMC. It will be important to study the preliminary analysis steps and their impact on the final extraction step in future work, especially if our formalism is to be adapted for use with ground-based experiments."
    },
    "1107/1107.3590.txt": {
        "abstract": "From 2005 to 2009, 25 observations of Cyg X-1 were performed with Suzaku, achieving a total exposure of 446 ks. In all observations, the source was found in the low/hard state, while the 1.5--12.0 keV count rate of the All-Sky Monitor onboard RXTE varied by a factor of $\\sim$ 3. In each observation, the 10--60 keV HXD-PIN spectrum and the 60--400 keV HXD-GSO spectrum were  fitted successfully by a thermal Comptonization model plus reflection by a thick neutral material. As the soft X-ray intensity increased, the Compton $y$-parameter was found to decrease from 1.0 to 0.6, while the solid angle of reflection to increase by $\\sim 30\\%$. Also conducted was timing analysis over a frequency range of 10$^{-3}$--10 Hz. As the source became brighter in soft X-rays, the characteristic frequency of hard X-ray variation increased from 0.03 to 0.3 Hz, while the fractional hard X-ray variation integrated over 10$^{-3}$--10$^{-2}$  Hz decreased by a factor of $\\sim$ 5. The signals in the 60--200 keV band were generally found to vary on shorter time scales than those in the 10--60 keV band. These spectral and timing results can be consistently interpreted by presuming that increases in the mass accretion rate cause the Comptonizing hot corona to shrink, while the optically-thick disk to intrude deeper therein. ",
        "introduction": "Black hole binaries (BHBs) are known to reside mainly in either of two spectral states, the low/hard state (LHS) and the high/soft state (HSS), which appear when mass accretion rate is lower and higher, respectively (McClintock \\& Remillard 2006). In the HSS, the accreting matter is considered to release its energy predominantly in soft X-rays as a multi-temperature disk emission (Mitsuda et al. 1984, Makishima et al, 1986) from an optically thick and geometrically thin accretion disk, called standard disk (Shakura \\& Sunyaev 1973). In the LHS, in contrast, the spectrum becomes much harder, and is characterized by a flat power-law with a photon index of $\\Gamma \\sim 1.5$ and a high energy cutoff at $\\sim$100 keV . This has been interpreted as inverse Comptonization of some soft seed photons by hot Maxwellian electrons (e.g., Shapiro et al. 1976; Sunyaev \\& Tr\\\"{u}mper 1979) in a geomatrically thick and optically thin accretion flow, or ``corona\", in which advection may be operating (e.g., Ichimaru 1977; Narayan \\& Yi 1995).  A scenario of the LHS,  called  ``disk-corona\" view (e.g., Liang et al. 1977), assumes that the system still harbors a standard disk, which provides the seed photons for Comptonization. Although the ``disk-corona\" model became a standard interpretation of the LHS of BHBs, the overall geometry of the disk and the corona remained poorly understood, including in particular, whether the geometrically thin disk extends to the last stable orbit (e.g., Miller et al. 2006a,b; Rykoff et al. 2007) or is truncated at some larger radii (e.g., Ebisawa et al. 1996; Poutanen et al. 1997; Gierlinski et al. 2008; Nowak et al. 2011). Equally unsettled is the interpretation of the rapid variation (e.g., Miyamoto et al. 1989) in terms of the disk-corona view. It was not clear, either, how these components evolve as the system approaches the HSS through the LHS. These limitations have been caused mainly by observational difficulties, to obtain high-quality spectra on short time scales, and over an energy range wide enough to determine simultaneously and precisely both the disk emission around $\\sim$1 keV and hard continuum shapes including the high energy cutoff energy of $\\sim$ 100 keV.  In 2005 October, a broad-band Suzaku observation of Cyg X-1 in the LHS was conducted for an exposure of 17 ks (Makishima et al. 2008; hereafter Paper I). The obtained 0.7--300 keV spectrum has been reproduced successfully by assuming that the disk is truncated at  $\\gtrsim 15 R_{\\rm G}$, and it provides the corona with the seed photons. Here, $R_{\\rm G}=G M_{\\rm BH}/c^2$ is the gravitational radius, with $G$ the gravitational constant, $M_{\\rm{BH}}$ the black hole mass, and $c$ the speed of light. The data also required a reflection hump with a solid angle of $\\Omega/2\\pi \\sim 0.4$ and a mildly broadened iron line with a breadth of $\\sigma \\sim 1$ keV; this is consistent with the disk truncation at $\\sim$ 15 $R\\rm{_G}$. The corona was inferred to be highly inhomogenious, expressed by a two-zone approximation invoking optical depths of $\\sim 0.4$ and $\\sim 1.5$, and the same electron temperature of $\\sim$ 100 keV. In addition, variations on a time scale of 1 s were studied in terms of Comptonized emission using intensity-sorted spectral analysis. It was then indicated that the X-ray intensity flares up, on time scales of $\\sim$ 1 s, when a larger fraction of the disk is covered by the corona, while the disk itself stays almost constant. In other words, the rapid X-ray flickering of Cyg X-1 in the LHS is primarily generated by fluctuations of the number of disk photons intercepted by the corona. This view is consistent with the one, based on the almost energy-independent fractional rms variation spectrum, that the variation in the LHS is generated by fluctuations in the seed photon input (Gierli\\'{n}ski \\& Zdziarski 2005; Gierli\\'{n}ski et al. 2010). Actually, figure 8(b) of Paper I  (high-phase to low-phase spectral ratio), which is almost equivalent to such rms spectra, is rather flat, except in $\\lesssim$ 1 keV where the ratio decreases due to the dominance of the disk emission (Yamada 2011). Thus, the results of Paper I reinforced the disk-corona picture, and quantified its details.  Given this snap-shot result, our next task is to clarify long-term evolution of the disk-corona configuration, which is a key to understanding the transition from the LHS to the HSS. For this purpose, we analyzed 25 Suzaku data sets of Cyg X-1 acquired over 2005--2009. Among the observations, the spectral and timing properties in hard X-rays (10--200 keV) showed significant changes, even though the source remained in the LHS throughout. In the present paper, we limit our analysis to the data obtained with the Hard X-ray Detector (HXD; Takahashi et al. 2007; Kokubun et al. 2007). Results with the X-ray Imaging Spectrometer (XIS; Koyama et al. 2007) will be reported elsewhere, after the data are fully corrected for pile up effects.  %%section 2 ",
        "conclusions": "Analyzing  the Suzaku HXD data of  Cyg X-1 acquired on 25 observations over 2005 October through 2009 December, we have obtained the following results. %% factor \u0082\u00a2\u0082\u00ad\u0082\u00c2\u0081A\u0082\u00c6\u0082\u00a2\u0082\u00a4\u0090\u0094\u008e\u009a\u0082\u00cd\u0081A\u0096{\u0095\u00b6\u0082\u00e2\u0090}\u0082\u00c6\u0093\u009d\u0088\u00ea\u0082\u00b5\u0082\u00c4\u0082\u00ad\u0082\u00be\u0082\u00b3\u0082\u00a2\u0081B %% \u0088\u00c8\u0089\u00ba\u0082\u00cc\u008e\u0096\u008e\u00c0\u0082\u00cd\u0081A\u008d\u00a1\u0089\u00f1\u0082\u00cc\u008a\u00cf\u0091\u00aa\u0082\u00c9\u0082\u00c2\u0082\u00a2\u0082\u00c4\u0082\u00cc\u0082\u00dd\u008c\u00be\u0082\u00a6\u0082\u00e9\u0082\u00b1\u0082\u00c6\u0082\u00c5\u0081A\u0082\u00dc\u0082\u00be\u0088\u00ea\u0094\u00ca\u0098_ %% \u0082\u00c5\u0082\u00cd\u0082\u00c8\u0082\u00a2\u0082\u00cc\u0082\u00c5\u0081A\u008c\u00bb\u008d\u00dd\u008c`\u0082\u00cd\u008eg\u0082\u00ed\u0082\u00c8\u0082\u00a2\u0082\u00e6\u0082\u00a4\u0082\u00c9\u0082\u00b5\u0082\u00dc\u0082\u00b5\u0082\u00bd\u0081B  \\begin{enumerate} \\item Cyg X-1 was  in the LHS on these occasions, accompanied by a factor 3 changes in $C_{\\rm{ASM}}$. The estimated 0.1--400 keV unabsorbed flux  varied over (4--8) $\\times$ 10$^{-8}$ erg cm$^{-2}$ s$^{-1}$.  \\item The 10--400 keV HXD spectra were all reproduced  successfully with a single Comptonization model, considering reflection from a cold matter.  \\item As $C_{\\rm{ASM}}$ became higher, the HXD flux and the reflection solid angle increased, while the Compton $y$-parameter decreased.  \\item When $C_{\\rm{ASM}}$ increased by a factor of $\\sim$ 3, the PSD break frequency became higher by an order of magnitude, and the power integrated over 10$^{-3}$--10$^{-2}$ Hz diminished by a factor of $\\sim$ 5.  \\item In the ACF analysis, the 60--200 keV signals always showed shorter correlation lengths than the 10--60 keV ones.  \\item The CCF revealed skewness so that the 60-200 keV signals were delayed from those in the  10-60 keV, and the effects became more prominent towards  higher mass accretion rate.  \\end{enumerate}  In order to consistently explain items 3 and 4 above, we have developed an interpretation that  a gradual increase in the mass accretion  rate causes the corona to diminish in size, while the cool disk to intrude more deeply into the corona. Property 5 and 6 listed above may be explained by considering a radial gradient in the corona, so that harder photons are produced in closer vicinity of the black hole."
    },
    "1107/1107.0820_arXiv.txt": {
        "abstract": "The observation of the Sun with a pinhole in order to check the first and the second Kepler's law has been realized for high school students. The settling of the experiment in order to minimize the errors of measure has permitted to verify the Rayleigh formula for obtaining sharpest images given the diameter of the pinhole. ",
        "introduction": "Aristoteles already observed that far from an hole of whatever shape the image of the Sun above a screen was always a circle. For him it remained an unsolved problem. Leonardo da Vinci also worked on this topic, exploiting the dark camera principle mainly for drawing. Kepler in 1600 invented the term {\\em dark camera} and used it for studying the partial phases of the total eclipse of the sun in Graz.\\cite{lom} Pinhole photography is nowadays a fine hobby, with several fans, and a conspicous literature about it was produced up to now. The formula of Rayleigh describes the relationship between the pinhole diameter and optimum focal distance,\\footnote{http://home.sol.no/~gjon/pinhole.htm} being $\\lambda=5.5 \\cdot 10^{-7}m$ the wavelength of the light:  \\be d=1.9\\sqrt{\\lambda f}. \\ee  We have found empirically this optimum value following a logical criterium, very instructive to be repeated. ",
        "conclusions": "This experiment has been improved either empirically putting all the apparatus groundbased, as theoretically driving the sampling of measures from small to greater focal length in order to diminish the systematic error. We can conclude that the precision of this method is $\\pm ~~30 ~~{\\rm arcsec}$, once eliminated the systematic error with the correction factor $r$. The factor $r$ is to be calculated (or interpolate) for each pinhole diameter and for each focal length, knowing the true value of the Sun diameter (or of another standard source). The precision $\\pm ~~30 ~~{\\rm arcsec}$ of this method is confirmed by the successfull observation of the solar spots. That precision is enough for detecting the Earth orbit ellipticity, and it is the best one in Astronomy without using the telescope. That procedures followed for the visible light are the same in the case of X-rays. Pinhole devices have been used in X-ray astronomy for the imaging in such waveband, and that experience can straighforward introduce such techniques to the students."
    },
    "1107/1107.2963_arXiv.txt": {
        "abstract": " ",
        "introduction": ". \\hspace{0.5cm}  Supersymmetric Grand Unification based on the SO(10) gauge group has received well deserved attention over the last 3 decades. Models proposed fall into two broad classes : those that preserve R-parity down to low energies \\cite{aulmoh,ckn,abmsv,ag1,ag2,rparso10,blmdm} using as Higgs the special representations (${\\bf{\\oot}}$) of SO(10) that contain R-parity even SM singlets and another large class of SO(10)   R-parity violating models\\cite{nRGUTs} that attempt to construct viable models using sets of small SO(10) representations even after sacrificing the vital distinction provided by  R-parity between matter and Higgs multiplets.  Besides structural attractions, such as the automatic inclusion of the conjugate neutrino fields necessary for neutrino mass, SO(10) GUTs offer  a number of other  natural features.  Among these are  third generation yukawa unification\\cite{3Gyukuni,hallrathempf}, automatic embedding of minimal supersymmetric Left-Right models, natural R-parity preservation\\cite{rparso10} down to the weak scale and consequently  natural LSP WIMP dark matter, economic and explicitly soluble  symmetry breaking   at the GUT scale\\cite{abmsv}, explicitly calculable superheavy spectra\\cite{ag1,bmsv,fuku04,ag2}, interesting gauge unification threshold effects\\cite{ag2,gmblm,blmdm,nmsgutI,nmsgutII}  which can lead to a  natural elevation of the unification scale to near the Planck scale\\cite{nmsgutI}, GUT scale threshold corrections to the QCD coupling $\\alpha_3(M_Z)$ of the required\\cite{langpol} sign and size\\cite{precthresh} and a deep interplay between the scales of Baryon and Lepton number violation as suggested by   the neutrino oscillation measurements and the   seesaw formulae connecting neutrino masses  to the B-L breaking scale.  The   fascination of the MSSM RG flow at large $\\tan\\beta$ stems from  the  tendency of third generation yukawa couplings  to converge, at the  MSSM unification scale \\cite{3Gyukuni,hallrathempf}, in a  manner   reminiscent of gauge unification in  the MSSM  RG flow\\cite{marsenj,amaldimssm}. For suitably large $\\tan\\beta$ and for close to central input values of SM fermion couplings at the Susy breaking scale  $ M_S \\sim M_Z$, third generation yukawas actually almost coincide at $M_X$.  On the other hand, in SO(10) theories with only the simplest possible fermion mass giving  Higgs content (a single \\textbf{10}-plet), when all the complications of threshold effects at $M_X\\sim 10^{16} GeV$(not to speak of those at   seesaw scales $M_{\\bar\\nu}\\sim 10^7-10^{12}\\, GeV$) are ignored, one does expect to generate boundary conditions for the gauge and yukawa couplings that are unified gauge group wise and (third generation) flavour wise.  However, fitting the rest of the known fermion data (15 more parameters) definitely requires other Fermion  Higgs multiplets (more \\textbf{10}-plets, $\\mathbf{120,\\oot}$ s etc). A principled position (\\emph{monoHiggism}?)  with regard to the choice of Higgs multiplets responsible for fermion mass (FM Higgs) is to accept \\emph{only one} of each irrep present in the conjugate of the direct product   of fermion representations. This principle may be motivated by regarding the different Higgs representations as characteristic ``FM channels'' through which the fermion mass (FM) is transmitted  in structurally distinguishable ways. For example the Georgi-Jarlskog mechanism distinguishes the \\textbf{45} plet Higgs in SU(5) ($\\boot$  in SO(10)) from the $\\mathbf{5+{\\bar 5}}$ (\\textbf{10} in SO(10)) due to their ability to explain the quark-lepton mass relations in the second and third generations respectively. Similarly the $\\boot$ in SO(10) is peculiarly suitable for implementing the Type I and Type II seesaw mechanisms for neutrino mass. If one duplicates  the   Higgs multiplets transforming as  the \\emph{same} gauge group representation,  for example by taking multiple \\textbf{10}-plets in SO(10), then one abandons  the quest for a structural explanation of the pattern of fermion masses in favour of ``just so'' solutions.   Moreover the many Higgs doublets typically present in the theory mix with the \\textbf{10}-plet derived Higgs doublets  to make up the light doublets of the effective MSSM. Both these features dilute the expectation of exact third generation  yukawa unification.  The technical difficulties of $b-\\tau$ unification  are attendant upon the fact that the GUT scale values of the third generation Yukawa couplings, in sharp contrast to the three gauge couplings, are highly infrared sensitive(see \\cite{wells} for a cogent discussion).  The accuracy of unification of third generation Yukawa couplings after the MSSM RG flow    up to the 1-loop GUT scale  $M_X^0(\\sim 10^{16.25}$ GeV) is strongly dependent on the low energy values of the yukawa couplings in the MSSM and thus requires large $\\tan\\beta$. It is also strongly dependent on the precise values of $y_b^{MSSM}(M_Z)$ and $y_t^{MSSM}(M_Z)$(both of which can receive significant Susy threshold corrections) besides $\\alpha_3(M_Z)$ (which needs  significant GUT threshold corrections to be consistent with experiment\\cite{langpol}). The large threshold corrections \\cite{hallrathempf,hitanbetacorr} to the relations between Standard Model(SM) and MSSM Yukawa couplings for down type quarks precisely in the large $\\tan\\beta$ MSSM variants that are in focus in a SO(10) context, combined with infrared sensitivity of yukawa unification  casts  doubt on the naturalness of exact $\\mathbf{t-b-\\tau}$ unification. As  these yukawa couplings  have been measured more precisely over the same period as the investigation of   third generation yukawa unification in Susy SO(10) models,  the strain of  these models has increased with time, and various just so mechanisms are invoked to preserve the exact third generation unification. Yet it would be strange if Susy SO(10) had no mechanisms to accommodate not only this difficulty but even the full complexity of the fermion spectrum.   Many  detailed  investigations\\cite{3Gyukuni,hallrathempf} of the yukawa unification    problem since the early days have focussed on models based on small Higgs representations($\\mathbf{16,10,45}$) and  seek  to understand only the third generation fermion masses. These studies   assume  that no Higgs Multiplet other than the \\textbf{10}-plet could contribute significantly to the third generation masses and use that to justify the assumption that high scale threshold corrections to the yukawa couplings are negligible. In the few studies  that   attempt to tackle   the full fermion spectrum, recourse is    had to a   subset of the many  possible non-renormalizable operators that contribute to fermion yukawa couplings in the effective theory below the unification scale.  Since the first two generation yukawas are so small relative to the third generation  it is possible   to build models using  small non-renormalizable  contributions that allow one to fit three generation charged  fermion  data\\cite{andersonraby,lucasraby} and even keep the operator dimension five contributions to baryon violation within or near to experimental limits\\cite{lucasraby}.  When other aspects, such as GUT scale spontaneous symmetry breaking, are to be addressed multiple Higgs \\textbf{10,45,16}-plets are often introduced\\cite{nRGUTs}. These then require  ad-hoc discrete symmetries to replace the    R/matter parity that is   broken in such theories. These symmetries also serve to prevent  interference between the assigned functions of the multiple copies of low dimensional Higgs representation introduced (along with appropriately limited non-renormalizable interactions that in effect play the role of the higher dimensional Higgs irreps that were disallowed at the outset). With  stringent unification  constraints at GUT scales and large threshold corrections at the Susy breaking scale only narrow regions of the Susy-GUT parameter space remain viable. Moreover specification  of the soft Susy spectrum and couplings to ensure    simultaneous  cancellations to keep the threshold corrections to  $m_b^{MSSM}(M_S)$ small and the $b\\rightarrow s\\gamma$ rate consistent with current limits \\cite{raby,baer,wells} is needed.  In general  such investigations have focussed on evaluating various    non minimal effects such as non-universal gaugino masses, D-term contributions to scalar non universality etc, operating in specific (tight) corners of the soft susy parameter space, to evade the conflicting demands  of  infra red sensitivity, large infrared threshold corrections and accurate yukawa unification. A much smaller set of papers \\cite{aulmoh,ckn,rparso10,ag1,abmsv,gmblm,blmdm,bert3,nmsgutI} focusses on renormalizable models without ad-hoc symmetries  and confronts the full complexity of the spontaeous symmetry breaking problem conjointly with fitting the fermion spectra \\cite{nmsgutI,nmsgutII}.  Renormalizable    SO(10) Susy GUTs\\cite{aulmoh,ckn,abmsv} employ  not only the {\\bf{10}}-plet but also the other two SO(10)-allowed fermion mass (FM) Higgs  representations i.e the $\\mathbf{\\oot}\\,\\,\\mbox{and}\\,\\,{\\bf{120}}$. Renormalizability of the GUT model is strictly maintained throughout  as a criterion for the structure of the superpotential. As a result these theories have a very parameter-economical structure\\cite{aulmoh,ckn,abmsv,nmsgutI}. In 1992, just as yukawa unification came into focus it was proposed\\cite{babmoh} that use of only the $\\mathbf{10-\\oot}$ FM Higgs might be sufficient to account for the entire fermion mass spectrum \\emph{including} the neutrino masses and mixings in terms of their (15)  yukawa coupling parameters. This obviously attractive possibility was investigated as a \\emph{generic} possibility in SO(10) (i.e without using the formulae in terms of the fundamental GUT parameters but only having the symmetries and (naive\\cite{pinmsgut})parameter counting   implied by their FM Higgs  yukawa  structure). Although initially unsuccessful, increasingly refined studies based upon the increasingly well defined neutrino mass data ultimately showed\\cite{TypeIfits,TypeIIfits} that the generic  $\\mathbf{10}-\\boot$ parametrizations permitted accurate fits of all the known fermion mass data  using \\textbf{10}-plet domination of third generation mass terms combined with Georgi-Jarlskog mechanism for the second generation and   either  the Type I or Type II seesaw  mechanisms or a combination thereof to fit the neutrino data. Perforce  these generic models always assumed  freedom to dial the relative strength of the two seesaw mechanisms.These accurate fits were found even while assuming that the precisely known SM yukawa coupling uncertainties  persevered unchanged over the  RG flow  up to the GUT scales. In the light of the drastic susy threshold corrections at large $\\tan\\beta$ this assumption was  naive and peculiarly uninformed of the attention long paid\\cite{hallrathempf,hitanbetacorr} to the role of these corrections in the other strand of SO(10) GUTs.  The appearance of \\cite{antuschspinrath,rosserna} however reminded us of this essential feature    and we have thenceforth adopted the much more realistic estimates of theoretical and experimental uncertainties advocated by \\cite{antuschspinrath} in our subsequent work\\cite{nmsgutII}.  Just as the  intriguing results on the successful implementation of the Babu-Mohapatra proposal in generic SO(10) GUTs  emerged it was also shown\\cite{gmblm,blmdm,bert3,pinmsgut} that : (i) The use of MSGUT specific fermion mass formulae implied that neither seesaw mechanism could account for neutrino masses while also fitting the charged fermion data. (ii) Conclusions derived using generic formulae could not be transferred to  concrete and specific GUTs because of intrinsic difficulties in untangling highly nonlinear constraints placed by the underlying fundamental GUT on the generic parameters. (iii) In a specific scenario\\cite{blmdm} inclusion of the remaining allowed  FM Higgs (i.e the ${\\bf{120}}$-plet) could permit the generation of large enough Type I seesaw masses due to suppressed $\\mathbf{\\oot}$ couplings and charged fermion mass fits due to the $\\mathbf{10-120} $ combination.  The $\\mathbf{10-120} $ FM Higgs combination proposed by us to tackle the charged fermion fits  (with the $\\boot$ couplings too small to affect any but the first generation masses) was seemingly  shown\\cite{grlav} to face a generic    difficulty in providing large enough strange and down quark masses.  This was only to be expected in any scenario where the Georgi-Jarlskog  or other mechanism  to  distinguish  second generation  down type yukawas is  not implemented. This difficulty appears in a quite different light when one considers\\cite{nmsgutII} that   Susy threshold corrections at large $\\tan\\beta$ tend \\cite{hallrathempf,hitanbetacorr} to drastically modify  the effective value of the down and strange quark yukawas. Thus we proposed\\cite{nmsgutII} that the excellent fits obtainable  for the other (16) fermion data (besides $y_{d,s}$) justified searching for the susy threshold corrections to implement the required corrections to lower $y_{d,s}$  and  this could be achieved not only with free soft susy parameters at $M_S$ (as in version 1 of this paper\\cite{nmsgutII})  but even with just 5 GUT compatible $N=1$ supergravity(SUGRY), Non Universal Higgs Masses(NUHM), type soft parameters $ ( m_{ \\frac{1}{2}},m_0,A_0,m^2_{H,\\bar{H}})$ at $ m_Z$ specified at $M_X$\\cite{nmsgutII} and two more $(|\\mu|, B\\sim m_A^2)$ determined by electro-weak spontaneous symmetry breaking  and run back up to $ M_X$ to provide 7 parameters at $ M_X$ coding all the SUGRY-NUHM information.  These  developments have put detailed consideration of the role of ${\\mathbf{120-\\oot}}$ FM Higgs representations firmly on the agenda of Susy SO(10) yukawa unification. In this paper we argue that the assumption that third generation fermion yukawas are protected from large GUT scale threshold corrections associated with non-\\textbf{10} -plet FM Higgs irreps is facile. We show that it is belied in the class of Susy SO(10) theories\\cite{aulmoh,ckn,abmsv,blmdm,nmsgutI,nmsgutII} which actually face up to the task of accounting for \\emph{all} the available fermion mass data in a fully specified model  without invoking uncontrollable higher dimensional operators or ad-hoc symmetries and rely solely on SO(10) gauge symmetry, supersymmetry, and structural(parameter counting) minimality  as guiding principles. In particular the NMSGUT with a large $\\mathbf{120}$-plet coupling (the $\\mathbf{120}$-plet yukawa  is antisymmetric and hence has two eigenvalues that are equal in magnitude  and one that is zero) \\emph{requires} evaluation of the GUT scale threshold corrections to the fermion Yukawa couplings. The large number of fields,  \\textbf{120}-plet couplings  comparable to the \\textbf{10}-plet couplings,   and the fact that these threshold corrections arise from chiral supermultiplet  wave function normalization(so that any field that couples to a matter field or a Higgs doublet can run in the self energy loop) raises the possibility that these corrections may be far from negligible\\cite{dixitsher}. In this paper we show that this is in fact the case.  The particular importance of the wave function  corrections for the fermion data fitting program is that they can relax the  stringent constraint $y_b-y_\\tau\\simeq y_s-y_\\mu$ that we found \\cite{core,msgreb} operative at $M_X$ in SO(10) models with a \\textbf{10-120} FM Higgs system. Thus at least any model that employs the $\\mathbf{120}$-plet must be sufficiently specified to allow the calculation of these very significant corrections which can drastically change the yukawa couplings at $M_X$. Such a specification involves   solution of the spontaneous symmetry breaking, calculation of the GUT spectrum and couplings and a RG analysis of threshold effects based thereon. So far, to our knowledge, such a complete calculation is possible  only for the SO(10)  NMSGUT \\cite{blmdm,nmsgutI,nmsgutII}for which we give the dominant 1-loop contributions( certain sub-dominant mixing effects are postponed to a future calculation). The calculation of the (wave function) renormalization that gives rise to the threshold effects modifying the matter fermion Yukawa couplings highlights the issue of perturbative consistency. One finds that while the corrections to the matter fermion lines of the mass generating Higgs-Yukawa vertices are no more than 25\\%, the wave function corrections($\\Delta_{H,\\bar H}$) to the (light) MSSM Higgs lines can be much larger.   The most serendipitous scenario would be if searches that restricted these corrections to be less than 1 were able to find fits that are also consistent with Baryon decay limits. This is not what we have observed in our searches so far : accurate fits with $|\\Delta_{H,\\bar H}| <1$ are obstructed in achieving small values ($<< 10^{-18} GeV^{-1} $) of the Baryon violating dimension 5 operator Wilson coefficients and hence give baryon lifetimes $\\sim 10^{27}$ years. On the other hand  if one allows large values of $\\Delta_{H,\\bar H}$ the searches achieve lifetimes that are a million times or more  larger. Thus although threshold corrections to gauge couplings turned out to be mild due to cancellations \\cite{ag2,blmdm} in spite of early alarms\\cite{dixitsher} the wave function renormalization effects at one loop are indeed  large. Investigation of higher loop contributions is thus be called for to see if the perturbation theory is sensible. In this paper we take the view that it is necessary to pursue the investigation of the realistic features of the SO(10) while including the large 1-loop threshold effects even though full perturbative convergence may take long to prove( or indeed may never be possible : as for instance in the most precise known theory QED).  In Section \\textbf{2} we briefly review the structure of the NMSGUT\\cite{blmdm,nmsgutI,nmsgutII} to establish the notation  for presentation of our results  on    threshold effects in Section \\textbf{3} and   Appendix \\textbf{A}. In Section \\textbf{4} we present illustrative examples to underline the significance of the GUT scale threshold effects and the need to include them. In Section \\textbf{5} we present  examples of improved fits(specially with respect to acceptable $d=5$ operator Baryon violation rates).  In Section \\textbf{6}   we discuss the broad features of the emerging phenomenology of the NMSGUT. In Section \\textbf{7} we summarize   our conclusions  and discuss the   various   improvements in the fitting, RG flows and searches   and    that are called for. Appendix \\textbf{A} contains details of the calculation of  threshold effects  at $M_X$.  Appendix \\textbf{B} contains a discussion of the two loop RGE flow at large $A_0,|M^2_{H,\\bar H}|$ values which results in the unconventional susy spectra with normal s-hierarchy and gaugino masses   not even close to the $1:2:7$ ratio found at one loop in  Susy GUTs with universal gaugino masses. Appendix \\textbf{C} contains tables of parameter values for   additional example fits to provide a wider view of the possibilities and as inputs for phenomenological explorations of the viability of our results. . ",
        "conclusions": "This paper is the third of  a series \\cite{nmsgutI,nmsgutII} developing  the NMSO(10)GUT as a possibly  viable and complete theory of particle physics.  The emphasis is to develop the theory on the same lines and as explicitly as the the Standard Model. We have shown that the theory is sufficiently simple as to allow explicit calculation of the spontaneous symmetry breaking, mass spectra and eigenstates  and allows a computation of the RG flow in terms of the fundamental GUT parameters to the point where one can attempt to actually fit the low energy data, i.e the SM parameters together with the neutrino mixing data,  in its entirety. In carrying out this project the model meets a major obstacle in its inability to fit the unmodified down and strange quark yukawas in the MSSM (renormalized up to $M_X$), which it overcomes by staking its viability on the operation of large $\\tan\\beta$ driven threshold corrections -providentially known to be operative and important in this context- which lower these yukawas from their SM values by a factor of about 5. As a result the pattern- although unfortunately not the scale- of the Susy breaking parameters tends to become fixed by the need to preserve, or indeed somewhat raise, the b quark yukawa while lowering the  $d,s$ quark yukawas when crossing the SM-MSSM threshold(s) . Thus the major unknown for the model, indeed for the whole field of Susy, remains the mass of at least one of the light susy particles. Due to the tight interrelationship of the susy spectra which have been generated from just 5 soft susy parameters at $M_X$ and subjected to multiple stringent  demands such as EW symmetry breaking, yukawa modification, Baryon decay suppression, LSP consistency and so on, it seems very plausible that input of any one of the susy particle masses would   generate a prediction of all the susy  masses in the context of a fit that was pinned to yield that particular mass value along with satisfying the various requirements we have mentioned. In the absence of Susy discovery data and the difficulty of extending the mass ranges that can actually be claimed to have been excluded without egregious assumptions, we can still call upon the the model to stand by it's claim to be a viable theory of Dark matter and yield its stable Bino LSP in the $5-150$ GeV mass range preferred by Cold Dark matter WIMP scenarios. This requirement is sufficient to pin the prediction for the Susy spectrum to a specific and completely novel and distinctive \\emph{type}, if not yet to specific masses for specific particles. The model thus predicts light gauginos and discoverable first or second generation right chiral sfermions below 2 TeV and invisible third generation sfermions. Thus we arrive at the attractively dangerous conclusion that : \\emph{If as all other GUT scenarios derive, the third sgeneration emerges lightest the NMSGUT may be taken to be falsified.}  This remarkable model has thus added yet more feathers to the already long  list of its attractions. Besides providing a natural and minimal context for the supersymmetric seesaw  mechanism and the implementation of R-parity as a part of the gauge structure of unification, the theory   successfully accounts for the entire available fermion mass data in terms of its own    parameters in way consistent with all known phenomenology.  It also yields insight into a necessary structure of it's susy breaking parameters  and yields a viable and completely natural candidate for the LSP combined with a surprising and novel candidate for the scalar NLSP which can lead to an effective DM scenario. Furthermore the theory naturally pushes the unification scale towards the Planck scale and allows suppression of the dangerous $d=5$ B violation operators. The conflation of the Planck scale and the unification scale goes a long way towards alleviating a perennial problem of  renormalizable SO(10) models \\cite{trmin,tas} namely divergent couplings in the ultraviolet. The unification scale or rather the Landau pole above it becomes a  physical cutoff beyond which the theory enters a strongly coupled phase together with gravity.    Till a susy breaking soft mass is pinned by experiment  the development of the NMSGUT will continue by   facing up to technical challenges that we have postponed in the first phase of the definition of the model as reported in a series of papers(\\cite{ag1,ag2,abmsv,gmblm,blmdm,precthresh,core,msgreb} and the papers of the current triplet: \\cite{nmsgutI,nmsgutII} and this paper). We conclude by itemizing the important issues which may materially alter the numbers obtained by our fitting program so far.  \\begin{itemize} \\item We have, following\\cite{piercebagger}, chosen $M_Z$ as the scale at which we match the SM to the MSSM. On the one hand this is well motivated since, as we have seen, the weak  gauginos, i.e the Bino and Wino tend to emerge as light as we allow them : which is as light as experiment permits i.e $\\sim M_Z$ for the Chargino but much lower for the Bino( the connection between the two masses and the limit $M_{{\\tilde W}^\\pm}$ however does not let the Bino mass descend below about 25 GeV). On the other hand we cannot deny that the model has its own little hierarchy problem with the $\\mu, A_0,M^2_H$ parameters all lying in the range of 10-100 TeV. Clearly the threshold effects due to the large differences of the Susy particle spectra from $M_Z$ will cause appreciable threshold corrections to the naive RG running which takes a single undifferentiated Susy breaking scale and identifies it arbitrarily with $M_Z$. Some development of the techniques for incorporating multiple thresholds between the SM and MSSM has already taken place\\cite{box} and should prove useful.  \\item We have calculated only tree  tree masses for the susy particles  in the theory (the Electro weak symmetry breaking and thus the  Higgs masses were however calculated using  1-loop effective potential\\cite{piercebagger,porod}. This is not a very good approximation in the MSSM with small $A_0$ and it may be worse in scenarios with large $A_0$ parameters. The large $A_0$ parameters lead to large trilinear couplings ($A=A_0 Y_f$)   only for the third sgeneration since the Yukawas $Y_f$ for the other generations are so small and may significantly modify   the third generation sfermion masses. The complete formulae for calculating these modifications-at least in the flavour diagonal case - have long been known\\cite{piercebagger}. They will be incorporated in the next version  of our search codes.  \\item We have not yet incorporated the three thresholds associated with the heavy right handed neutrino masses in the theory. Since one progressively introduces the neutrino Dirac couplings as one passes these thresholds going higher in energy, significant effects on the yukawa unification may be anticipated. These thresholds will also be significant when calculating the amount of lepton flavour violation introduced when integrating down from the high scale (driven by the non diagonality of the SO(10) yukawas required to account for the observed  quark flavour mixing). It is for this reason, and because the mass insertion formalism is ill adapted to securely  evaluate novel scenarios, that we have not generated tables of Lepton flavour violating mass insertions for comparison with the existing analyses on lepton flavour constraints\\cite{masvemp}. We note however that we  have calculated some of the common Susy sensitive quantities and found them to be compatible with existing limits. Nevertheless the incorporation of constraints on  our searches based upon electric dipole moments and  the strength of flavour violation operators involving the sfermions of the first two sgenerations are a priority issue for our program.  \\item It will be clear to the reader familiar with the nuances of the Susy Flavour problem that in terms of a Bottom-up approach our results suggest that a radical extension of the consistent and viable susy parameter space at large $\\tan\\beta, A_0,|mu|,|M^2_{,\\bar H}|$  and with $M_{\\tilde f_3}>>M_{\\tilde f_3}$  may be possible and the characteristics of such an extension as derived from the NMSO(10)GUT approach could be diametrically opposed to those from all previous GUTs. However that reader will also realize that the inversion of the standard inverted hierarchy of sfermion masses to a normal hierarchy may have drastic implications for the consistency of the theory with flavour violation constraints and EDM constraints for the first two generations (the heavy sthird is likely to pass such constraints -as already seen in the case of the $b\\rightarrow s\\gamma$ constraint : which is commonly a stringent one when the third sgeneration is light but is here totally insensitive see Table \\ref{LWENRGYCNSTRTS}). Prominent among these well known constraints are those on the electric dipole moments of the electron, muon,  neutron etc, the severely constrained Branching rations of the $\\Delta F=1$ decays of Mesons such as $B_d\\rightarrow \\mu\\mu$, and the limits on the supersymmetric contributions to $\\Delta F=2$ quantities like $\\epsilon_K ,\\Delta K$ etc. These contributions can now be calculated using codes\\cite{Susy-flavour} that do not make any simplifying approximations as done in the mass insertion method\\cite{gbgbrmas} which allow only order of magnitude estimates of individual terms but cannot be used to evaluate total contributions including the effect of cancellations, in unconventional scenarios like ours. It may well be that the parameter sets we have presented will fail these tests when they put to them. However we emphasize that, from the point of view of the this series of papers which seeks to pin down the viable points   of the 45 dimensional  parameter space SO(10) NMSGUT, such a challenge would be no different than the one that was posed by the onus to show compatibility with B violation rates:  which the theory overcame by novel use of expedients available to it in a way  that pointed out fresh approaches to hoary questions. In the same way it may be that when the requirement of the consistency of the \\emph{total} value of such parameters predicted by the NMSGUT is taken into consideration then again new viable parameter sets may emerge. In view of the length of the present paper as well as the considerable further effort required to answer definitively to these important challenges the flavour violation constraints will be taken up in the sequels. We urge the tolerant reader not to prejudge this vital issue but join us in reflecting on the behaviour of the  novel type of Susy  parameter sets suggested by us.  \\item In this calculation we  found parameter  sets leading to quasi-stable Baryons by a shotgun carrying brute search. However the characterization of the possible cancellations and the least constraining  versions compatible with experimental limits remains to be done. The search for fits with suppressed Baryon Number violation was carried out by simply limiting the size of the maximal element of the LLLL and RRRR operator coefficients with no heed paid to the detailed  effects of that coefficient on the decay rate in specific channels. A more sophisticated  (but hugely more computer intensive) way of doing this would be to calculate the Baryon violation rates in each channel at every iteration and limit the total lifetime. In principle one could hope to implement this given enough (super)computing power.  \\item  Due to the large amount of running time required to find an acceptable solution, we have only scratched the surface of the enormous parameter space and can by no means pronounce on the general structure of the solution space. It will take  long runs on a super computing cluster to develop a statistical picture of where the solutions tend to lie. We are now preparing for both the improvements mentioned above and the harnessing of a cluster  for the task.  \\item We have run down only the diagonalized yukawa couplings from $M_X^0$ to $M_Z$ along with the susy breaking parameters which are optimized to fit the eigenvalues of the SM yukawa couplings to the run down diagonal NMSGUT matter yukawas. A complete treatment would run down the full set of coupling matrices  obtained from the GUT and apply the large $\\tan\\beta$ driven Susy threshold corrections to the off diagonal couplings before matching the two sets : or at least the ``physical'' parameters (eigenvalues, mixing angles and phases) coded in the two pairs. The formalism for applying off diagonal corrections is still somewhat murky  so some theoretical progress towards setting out clear algorithms for including all off diagonal 1-loop effects and renormalizations is called for.  \\item We argued that the CCB and UFB minimae(that certainly exist at the large values of $|\\mu|,|A_0|$ central to the needs of the NMSGUT) need not destabilize the metastable standard vacua we have found and cited previous investigations of this issue\\cite{kuslangseg} which are specially encouraging in this regard. The fact remains, however,  that these decay rates must actually be estimated for each set of  parameters found. On the other hand once the requisite subroutines for calculating the vacuum tunneling rate are in hand one can add them to the search routines to filter the parameter sets.    \\item Threshold corrections play a central role  in our calculation and the large wavefunction dressing of the Higgs doublet lines that we find at one loop demand an investigation of the two loop effects to determine whether they are yet larger still. If so our model, and by implication realistic GUTs generally,  may survive only after the Higgs wavefunction dressing has somehow been re-summed to all orders. However we note that this growth of wave function dressing leads to reduction in the actual size of the SO(10) yukawas actually needed to fit the low energy fermion data by a factor of 10-100. Thus the contributions to the fermion lines are effectively degenerate and it is only the large number of heavy fields running in dressing of the light Higgs boson lines that leads to  such a large  wave function renormalization of the Higgs fields. It is possible that in case of further growth at the multiloop level the theory is being driven to a quantum fixed point dominated by the gauge coupling and therefore re-summable  using the the exact beta functions available for supersymmetric gauge theories\\cite{vainshif}. In any event we consider that our calculations show that GUTs aiming to be realistic have been ``up-ended'':  in the sense that any Grand unified model  with pretensions to realism must define itself explicitly enough to permit calculation of   threshold corrections using calculated superheavy field spectra or else risk being under suspicion of being an qualitative scenario that may be destroyed as soon as quantum effects are included.  \\item As Karl Popper emphasized  so insightfully, the virtue of a comprehensive scientific theory that accounts for all known data is that it accepts the challenge and charm of living dangerously and in constant confrontation with experimental data that may falsify it. It must face every new measurement that it claims to be able to account for with its fate hanging  in the balance. This indeed is what gives properly scientific models   their  peculiar power and utility: that speculations living at a safe distance from the cutting edge of Occam's razor rarely possess. The (N)MSO(10)GUT has braved multiple iterations of challenges over the three decades since it was proposed\\cite{aulmoh,ckn} and, so far,  has emerged stronger from every challenge,  defining the possible in its realm ever more clearly and distinctly. If it dies it will not have lived in vain. \\end{itemize} \\bea\\nonumber\\eea"
    },
    "1107/1107.0966_arXiv.txt": {
        "abstract": "We report the discovery and characterization of a power law correlation between the local surface densities of Spitzer-identified, dusty young stellar objects and the column density of gas (as traced by near-IR extinction) in eight molecular clouds within 1~kpc and with 100 or more known YSOs.  This correlation, which appears in data smoothed over size scales of $\\sim$1 pc,  varies in quality from cloud to cloud; those clouds with tight correlations, MonR2 and Ophiuchus, are fit with power laws of slope 2.67 and 1.87, respectively. The spread in the correlation is attributed  primarily to local gas disruption by stars that formed there or to the presence of very young sub-regions at the onset of star formation. We explore the ratio of the number of Class~II to Class~I sources, a proxy for the star formation age of  a region, as a function of gas column density; this analysis reveals a declining Class~II to Class~I ratio with increasing column density. We show that the observed star-gas correlation is consistent with a star formation law where the star formation rate per area varies with the gas column density squared. We also propose a simple picture of thermal fragmentation of dense gas in an isothermal, self-gravitating layer as an explanation for the power law. Finally, we briefly compare the star gas correlation and its implied star formation law with other recent proposed of star formation laws at similar and larger size scales from nearby star forming regions. ",
        "introduction": "}  It has been long recognized that the density of young stars varies by orders of magnitude between different star forming regions.  \\citet{herb62} compared the low density of young stars in the Taurus star forming region to the high density of young stars in the Orion nebula and speculated that these may form an evolutionary sequence, with the low density regions collapsing into denser structures.   More recently, ground-based near-IR surveys of molecular clouds \\citep[e.g.][]{lada91,stro93,carp00,porr03,lada03} identified embedded dense young stellar clusters and groups in molecular clouds as well as a more diffuse population of low mass stars, showing that molecular clouds contain both a dense ``clustered'' and a diffuse ``distributed'' population.  Although clusters could easily be identified in these surveys by searching for peaks in the surface density of sources, measurements of the number of distributed young stars in the clouds were dominated by uncertainties in the amount of contamination by field stars and the treatment of variable extinction through each cloud \\citep{carp00}.  This led to a number of questions regarding the number and origin of young low mass stars outside of clusters \\citep{lada91, li1997,carp00}.   What fraction of the cloud population is found outside of clusters? Are they born in clusters and dynamically ejected to their observed positions \\citep{evan09}?  Alternatively, if they are formed near their current positions, why do molecular clouds exhibit large variations in the density of young stars?  Are the variations in density related to variations in the Jeans length within clouds \\citep{teix06,gute09}?  Better observational characterization of the relationship between the spatial distribution of YSOs and their natal cloud material is needed to understand the nature of the distributed component, constrain models of star formation in molecular clouds, and ascertain the underlying physics that determine the density and rate of star formation \\citep[e.g.][] {bate05, krum07, myer09}.   Sensitive {\\it Spitzer} mid-IR and 24~$\\mu$m surveys of nearby young clusters embedded in molecular clouds have enabled complete characterization of the full range of stellar densities in a range of star foming environments \\citep[e.g.][]{alle07,evan09}.  Young stellar objects (YSOs) in the clouds are identified by the presence of IR-excess emission from circumstellar disks and/or envelopes.  Several groups have developed methods to use {\\it Spitzer}-determined YSO spectral energy distributions to characterize the sources as protostars (Class~0, I, or flat-spectrum) or pre-main sequence stars with disks (Class~II or transition disks) \\citep[e.g.][] {gute09, evan09, rebu10}.  By employing strategies to minimize the contamination by background galaxies with similar infrared colors \\citep[e.g.][]{harv06, gute08, gute09, evan09}, the {\\it Spitzer} surveys have yielded reliable, high spatial dynamic range distributions of YSOs in their natal clouds.  These observations show that the clusters are peaks of a more extended population of young stars which tend to follow the morphology of the filamentary molecular clouds \\citep{alle07,gute09,evan09}.  Thus, low density star forming regions are not limited to Taurus-like regions, but are also found in molecular clouds forming large clusters and high mass stars \\citep{alle07}.   The wide range in observed YSO surface densities provides an opportunity to study how the density of young stars varies with the observed star formation efficiency and the properties of the co-spatial molecular gas.  Recent work indicates that the star formation efficiency increases with the stellar density.  \\citet{evan09} showed that higher surface density YSO clusterings tend to exhibit higher star formation efficiency (30\\%) than lower density surroundings (3-6\\%).  In a {\\it Spitzer} survey of YSOs in the W5 molecular cloud, where a recent generation of star formation is proceeding on the peripheries of the evacuated cavities of adjacent HII regions, \\citet{koen08} found star-formation efficiencies of $>$10-17\\% for high surface density clusterings and 3\\% for lower density distributed regions, respectively.  The fraction of {\\it Spitzer} identified YSOs outside of clusters depends strongly on the adopted functional definition of stellar cluster, as there is no strong distinction in source spacings evident in the nearby YSO spacing data to suggest distinct ``clustered'' and ``distributed'' star-forming modes \\citep{bres10}.  The continuous nature of the overall YSO surface density distribution in nearby clouds motivates a different approach whereby the local surface density of YSOs and the local column density of gas are compared.  This approach does not require a functional definition of clustered or distributed star formation.   A similar approach has been used with success in observations of external galaxies.  \\citet{kenn98} outlines previously published results concerning the validity of the Schmidt Law of star formation in external galaxies, namely $\\Sigma_{SFR} = A \\Sigma_{gas}^N$, where the $\\Sigma$s refer to the observed surface densities of star formation rate and gas respectively \\citep{schm59}. For a limited range of densities and galaxy types these studies found $N$ to be between 1 and 2.  \\citet{kenn98}, using H$\\alpha$ emission as the star formation tracer, and CO and HI as the gas tracers for spiral galaxies, and Br$\\gamma$ emission and FIR dust emission as star and gas tracers respectively for infrared-selected starburst galaxies, show $N = 1.45 \\pm 0.15$.  Note that these observations measure the integrated star formation (as traced by recently formed OB stars) across multiple molecular clouds, hence the extragalactic laws should not be expected to apply directly to studies of low mass stars within individual molecular clouds in our galaxy.  That stated, \\citet{evan09} compare $\\Sigma_{SFR}$ against $\\Sigma_{gas}$ for the seven nearby clouds surveyed by the {\\it Spitzer} c2d Legacy survey against the Kennicutt-Schmidt Law, and find $\\Sigma_{SFR}$ to be $\\sim$20 times larger than predicted by the law for the measured $\\Sigma_{gas}$.   In this contribution, we perform a direct comparison between locally determined mass surface densities of YSOs and gas. This technique reveals a power law correlation between the YSO surface density and the gas surface density in all clouds, which we refer to hereafter as the star-gas correlation.   Furthermore, using the ratio of protostars to pre-main sequence stars as an evolutionary indicatory, deviations from that correlation can be interpreted as regions of more or less time evolution in the star-gas correlation. We interpret these by adopting a star formation law of power law form and calculating the integral of the star formation rate over time vs. the remaining gas density.  We present a simple analysis which shows that the higher density regions of a cloud produce stars which consume their local gas more rapidly, and the relation between log $\\Sigma_*$ vs. log $\\Sigma_{gas}$ increasingly deviates from a power law with time, particularly at high gas densities. Nevertheless, the observed correlation can be reasonably well explained by an underlying $\\Sigma_{SFR} = A \\Sigma_{gas}^{2}$ star formation law. We examine the observed deviations from the correlation and argue that they result from a combination of gas dispersion in rich clusters and non-coevality within the molecular clouds.   In Section~\\ref{obs}, we summarize the {\\it Spitzer} YSO lists compiled from prior observations, as well as the near-IR extinction maps (as a tracer for the molecular gas surface density).  In Section~\\ref{corr}, we demonstrate and characterize the observed correlations.  In Section~\\ref{discuss}, we present some analysis of the correlation and its implications including: a simple evolutionary model to demonstrate the robustness of an underlying star formation law of power law form under the effects of early gas depletion; a simple thermal fragmentation model that predicts the central values of the ranges of power law index and normalizations observed; a brief comparison to the Schmidt-Kennicutt star formation law for galaxies.  Finally, we summarize our findings in Section~\\ref{summary}. ",
        "conclusions": "}  The Class~II YSO evolutionary stage lasts between 0.5 and 5~Myr \\citep{hern08}, and the majority of the IR-excess sources included in this analysis are of that class.  Therefore, the fact that the above power law trend is observed at all suggests a close link between the YSO spatial distribution and the gas structure on the timescale of $\\sim$2~Myr for the $\\sim$1~pc size scales probed here.  While this is indeed a profound result, it is still in some sense unconstraining. The connection between the gas and the YSOs may in part result from the observed kinematic connection between dense cores and the surrounding molecular gas.  Specifically, millimeter-line observations show that the the velocity dispersion between the  cores and the surrounding, moderate density gas is small relative to the velocity dispersion inferred from the observed linewidth and the observed velocity gradients in the bulk gas motions \\citep[e.g.][]{wals04,wals07,adam06,kirk07,jorg08}. However, the YSOs and their progenitor gas could simply co-move in aggregate around the larger structure of the cloud, perhaps even dynamically interacting with other such aggregations, and still preserve the observed connection between the YSO and gas distribution. Thus, a dynamical picture for the cloud and the structures within the cloud is not ruled out by the star-gas correlation.  Under the simplifying assumption that there is relatively little motion between the YSOs or the gas, we present a simple model which adopts a star formation law where the star formation rate per area follows a power law of the local instantaneous gas surface densities.  This model demonstrates that the likely gas depletion that has occured to generate the observed star-gas correlation still leaves behind a power law of similar index to the primordial relationship. We then establish that the observed star-gas correlations and the star formation law have a power law index $\\sim$2. Next, we demonstrate that simple thermal fragmentation of an isothermal, self-gravitating, modulated layer of gas can produce a spatial distribution of gravitationally unstable fragments where the surface density of fragments scales as the gas density squared.  Finally, we briefly compare the reported star-gas correlation to similar recent analyses in the literature.   \\subsection{Comparing Star Formation Laws to the $\\Sigma_*$ vs. $\\Sigma_{gas}$ Relationship\\label{tommodel}}  As discussed in Section~\\ref{intro}, several empirical star formation laws have been proposed in the literature; these laws relate the inferred star formation rate per area to the current column density of gas, and typically they are derived over large ($>$100~pc) size scale measurements.  We examine here whether the observed $\\sim$1~pc scale $\\Sigma_*$~vs.~$\\Sigma_{gas}$ correlation reported in Section~\\ref{corr} is consistent with a simple star formation law where the star formation rate per area varies linearly or as a power of the current gas column density.  In order to consider the likelihood that the reported power law correlation in observed quantities, ie. the integrated number of YSOs and the remaining amount of gas, relates at all to an underlying star formation law of power law form, we must consider whether the more efficient creation of stars at higher gas column density significantly affects the underlying gas distribution and therefore results in an observed $\\Sigma_*$~vs.~$\\Sigma_{gas}$ distribution that deviates from the power law index of the star formation law itself.  To this end, we present a simple model of gas to stellar mass conversion.  Gas mass is accreted onto stars at a rate governed by an assumed star formation law with an efficiency that accounts for the ejection of gas mass from the system via protostellar outflows.  The nature of this model demands several underlying assumptions to both achieve adequate simplicity and offer a relatively robust examination of how the observables behave with time.  The first assumption is that the stars formed do not move significantly from their birth sites to within the typical resolution of our measurements ($\\sim$1~pc scale).  The observations that spatially proximate prestellar cores are likely formed with low relative velocities \\citep{wals07,muen07b} suggest that the typical velocity dispersion of of the YSOs that emerge from the cores is small ($\\sim$0.4~km~s$^{-1}$).  Furthermore, later in the star formation process, relatively uniform surface density distributions are observed in many embedded YSO groupings \\citep{gute09}, suggesting relatively little dynamical evolution.  The second assumption is that there is no molecular gas flowing into or within the molecular cloud at the $\\sim$1~pc scale; gas is only ejected from the system by the formation of stars, or converted into stellar mass.  This assumption is adopted for simplicity; it is demonstrably false at scales $<$1~pc, where systematic gas infall motions have been observed in cluster-forming clumps \\citep[e.g.][]{wm99,wals06}, and on larger scales, the global freefall time of a typical GMC is $\\sim$5~Myr, though turbulent support likely inhibits global collapse to some degree.  Finally, we assume that there is no effect on the gas distribution by large scale effects from phenomena such as feedback from nearby high mass stars, supernovae, or galactic motions.  These processes only affect the cloud on longer timescales ($>$10~Myr) than we are considering, with the exception of high mass star feedback that we ignore for simplicity.  We begin by adopting a general star formation law where the star formation rate per area shows a power law dependence on the gas column density:  \\begin{equation} \\frac{\\partial \\Sigma_{\\star}(x,y,t)}{\\partial t} = c k \\Sigma_{gas}(x,y,t)^{\\alpha} \\label{eqn:sflaw} \\end{equation}  \\noindent where $\\Sigma_{\\star}(x,y,t)$ is the mass surface density of YSOs as a function of position in the sky and time, $k$ is a constant that is defined in equation~\\ref{gasdepl} below, and $\\Sigma_{gas}(x,y,t)$ is the mass column density of molecular gas as a function of position and time.  We have assumed that the formation of a star depletes the cloud by $M_{\\star}/c$ where $M_{\\star}$ is the mass of the star created and $c$ is the mass conversion efficiency which takes into account the mass of gas ejected from the cloud by feedback during star formation ($M_{ejected} = M_{\\star}/c-M_{\\star}$).  Here we are assuming that the fraction of the mass in the outflow ejected from the cloud is constant throughout the cloud, but in fact the fraction of mass that escapes the cloud's gravitational potential may vary as a function of local gas column density and be affected further by the nearby gas configuration at $>$1~pc scales.  Furthermore, it is worth noting that in this formulation we are not considering individual stars, thus the entire analysis, including the value of $c$, is averaged over the stellar initial mass function (IMF), regardless of its form.  A pre-stellar core to star mass conversion efficiency that is constant over the stellar mass range is consistent with recent characterizations of the mass function of pre-stellar cores \\citep[e.g.][]{alve07,andr10}.  They find a value for this efficiency of $0.3 \\pm 0.1$.  However, this does not speak to the amount of unaccreted matter that is ejected from the cloud entirely; thus we consider 0.3 as a lower limit on the value of $c$.  If we assume that no additional gas is flowing into or within the cloud as the stars form, then the rate of depletion of the gas mass is given by:  \\begin{equation} \\frac{\\partial \\Sigma_{gas}(x,y,t)}{\\partial t} = - k  \\Sigma_{gas}(x,y,t)^{\\alpha} \\label{gasdepl} \\end{equation}  \\noindent In the case that $\\alpha \\ne 1$, the solution for $\\Sigma_{gas}(x,y,t)$, the gas density remaining in the cloud, is:  \\begin{equation} \\Sigma_{gas}(x,y,t) = \\Sigma_{gas}(x,y,0) (\\frac{t}{t_0}+1)^{\\beta} \\label{eqn:siggas} \\end{equation}  \\noindent where, $\\Sigma_{gas}(x,y,0)$ is the column density of gas at the onset of star formation, $\\beta$ is  a constant, and $t_0$ is the timescale for the gas to be depleted by a factor of $2^\\beta$.  Specifically:  \\begin{equation} \\beta = \\frac{1}{1-\\alpha} \\end{equation}  \\noindent and  \\begin{equation} t_0 = \\frac{1}{ k (\\alpha-1) } \\Sigma_{gas}(x,y,0)^{1-\\alpha} \\label{t0def} \\end{equation}  \\noindent Note that $t_0$ depends on the initial column density of gas for $\\alpha > 1$; in this case regions with higher column densities will be depleted more rapidly.  The amount of mass converted into stars is then the mass of the depleted gas times the mass conversion efficiency, $c$:  \\begin{equation} \\Sigma_*(x,y,t) = c [\\Sigma_{gas}(x,y,0) - \\Sigma_{gas}(x,y,t)] \\label{eqn:siggas_diff} \\end{equation}  \\noindent Substituting equation Eqn.~\\ref{eqn:siggas} into Eqn~\\ref{eqn:siggas_diff}, we obtain the relationship:  \\begin{equation} \\Sigma_{\\star}(x,y,t) = c\\Sigma_{gas}(x,y,0)[1-(\\frac{t}{t_0}+1)^{\\beta}]. \\label{fullsoln} \\end{equation}  In the case that $t << t_0$, we find that $\\Sigma_{\\star}(x,y,t)$ has the same power law dependence as the star formation rate:  \\begin{equation} \\Sigma_{\\star}(x,y,t) =  c \\Sigma_{gas}(x,y,0)^{\\alpha} k t \\label{simplesoln} \\end{equation}  \\noindent In other words, the surface density of formed stars is proportional to the surface density of initial gas to the $\\alpha$ power. In this limit, the observed gas density can be approximated as $\\Sigma_{gas}(x,y,t) = \\Sigma_{gas}(x,y,0)$ and the following approximation for the star formation rate can be used:   \\begin{equation} \\frac{\\partial \\Sigma_{\\star}(x,y,t)}{\\partial t} = \\frac{\\Sigma_{\\star}(x,y,t)}{t} =  ck \\Sigma_{gas}(x,y,t)^{\\alpha} \\label{alphagen} \\end{equation}   \\noindent In the case that the star formation rate is proportional to the gas density, i.e. $\\alpha = 1$ (even though our data cannot be reproduced by such a law, as we next show), we have a different solution to equation~\\ref{gasdepl} for the gas remaining, $\\Sigma_{gas}(x,y,t)$:  \\begin{equation} \\Sigma_{gas}(x,y,t) = \\Sigma_{gas}(x,y,0)e^{-kt} \\end{equation}  \\noindent and for $\\Sigma_{\\star}(x,y,t)$ we find the solution:  \\begin{equation} \\Sigma_{\\star}(x,y,t) = c \\Sigma_{gas}(x,y,0)[1-e^{-kt}] \\end{equation}  \\noindent which becomes, in the limit of $kt << 1$:  \\begin{equation} \\Sigma_{\\star}(x,y,t) = c \\Sigma_{gas}(x,y,0) k t \\label{alpha1} \\end{equation}  \\noindent Eqn.~\\ref{alpha1} thus simplifies to Eqn.~\\ref{simplesoln} for the $\\alpha=1$ case, under a similar assumption of small $t$.  At this point, we have argued that a power law observed in $\\Sigma_*$~vs.~$\\Sigma_{gas}$ is indicative of a star formation law of similar functional form.  Therefore, it is interesting to examine the $\\alpha =2$ case, as suggested by the analysis from Section~\\ref{corr}.  Combining Eqn.~\\ref{eqn:siggas}~and~\\ref{t0def} under the assumption that $\\alpha=2$ allows us to find a useful expression for the quantity $\\Sigma_{\\star}/\\Sigma_{gas}^2$ computed in Section~\\ref{corr}:   \\begin{equation} \\frac{\\Sigma_{\\star}}{\\Sigma_{gas}^2}(x,y,t) = ckt (1+\\frac{t}{t_0}) \\label{sigrat} \\end{equation}  \\noindent Thus for $\\alpha=2$, the histograms presented in Fig.~\\ref{histyso} can be thought of as histograms of normalized star counts for Class I and II YSOs vs.  $ckt_{SF}(x,y) (1+\\frac{t_{SF}}{t_0}(x,y))$ at the current mean age of the YSOs at a given position, $t=t_{SF}$.  Using this result and the observed evolutionary phase ratios of the YSO population, $N_{II}/N_I$~vs.~$\\Sigma_{gas}(x,y,t_{SF})$ (Fig.~\\ref{c2c1vgas}), we can constrain the constant $k$ and thus the gas depletion time $t_0(\\Sigma_{gas})$. Then we can directly compare our data to the gas depletion model.  Here we measure $t_{SF}(x,y)$ by computing the ratio of the surface density of stars formed divided by the star formation rate surface density.  We assume that the star formation rate surface density is constant in time and that it is approximated by the ratio of the surface density of Class~I objects over the lifetime of that evolutionary phase,\\ $\\Sigma_I/t_I$, where $t_I\\sim$0.5~Myr \\citep{evan09}.  Thus:  \\begin{equation} t_{SF}(x,y) = \\Sigma_*(x,y,t_{SF}) \\times t_I / \\Sigma_I(x,y) \\end{equation}  \\noindent $\\Sigma_*(x,y,t_{SF}) = \\Sigma_I(x,y)+\\Sigma_{II}(x,y,t_{SF})$ in this analysis, thus we can simplify to:  \\begin{equation} t_{SF}(x,y) = (\\frac{N_{II}(x,y,t_{SF})}{N_{I}(x,y)}+1) \\times t_I \\label{tsfdef} \\end{equation}  As shown in section~\\ref{corr}, we have computed $N_{II}/N_{I}$ for each cloud as a function of gas column density, noting that some clouds have a downward trend, while others are fairly stable over the whole range of density values. The more extreme variations occur in clouds with wider spreads in $\\Sigma_* / \\Sigma_{gas}^2$, thus we assume that those clouds have some contamination of lower gas column density regions by older YSOs (via their motion away from their birth sites) or that a large fraction of the primordial gas in those regions has been recently ejected by feedback from nearby, high mass stars. Given these complications, we have chosen to characterize the trend in $N_{II}/N_{I}$~vs.~$\\Sigma_{gas}$ from the relatively uniform MonR2 data, yielding:  \\begin{equation} \\frac{N_{II}}{N_{I}}(x,y,t_{SF}) = 25 \\times \\Sigma_{gas}(x,y,t_{SF})^{-0.4} \\label{n2n1} \\end{equation}  \\noindent The observed data are largely confined to the range $20 < \\Sigma_{gas}(x, y, t_{SF} ) < 300$~M$_{\\odot}$~pc$^{-2}$; evaluating eqn.~\\ref{n2n1} at these fiducial points yields $N_{II}/N_{I} = 7.5$~and~3, respectively, and $t_{SF} = 4.3$ and 1.8~Myr, respectively.  We now use these results  to set limits on the value of $k$, and thus $t_0$.  Since $t_{SF}$ and $t_0$ depend on quite different powers of $\\Sigma_{gas}$, the case of $t_{SF}<<t_0$ likely occurs at low initial gas column densities (this will be verified later), thus the higher value, $t_{SF}=4.3$~Myr, is applicable.  Assuming $c=0.3$ (the lower limit), then $2.3 \\times 10^{-4} < k < 3.9 \\times 10^{-3}$~Myr$^{-1}$~pc$^2$~$M_{\\odot}^{-1}$.  Therefore, at the corresponding value of $\\Sigma_{gas}$ (20~$M_{\\odot}$~pc$^{-2}$), where the depletion time is long and thus $\\Sigma_{gas}(x,y,t) = \\Sigma_{gas}(x,y,0)$, we find that the depletion time is $13 > t_0 > 220$~Myr\\footnote{Note that $t_0$ would increase (as k would decrease) linearly if $c$ is larger, and $c \\le 1$ by its definition above.}.  This range is 4 to 50 times larger than $t_{SF}=4.3$~Myr and thus consistent with the $t_{SF}<<t_0$ assumption. Assuming that the spread in the MonR2 $\\Sigma_*/\\Sigma_{gas}^2$ values are dominated by data uncertainty\\footnote{Our stellar surface densities are $n=11$ nearest neighbor densities, which are inherently uncertain at the 33\\% level \\citep{ch85}, and our gas surface densities are uncertain at 0.5-4~$A_V$, with higher uncertainty at higher surface densities.  These uncertainty sources result in $1\\sigma$ uncertainties of $\\sim$0.3~dex in $\\Sigma_*/\\Sigma_{gas}^2$ values, thus they are sufficient to account for the $\\pm0.6$~dex range in those values in MonR2, assuming that they are log-normally distributed and that our adopted limiting values isolate the central $\\pm2\\sigma$ range.}, we adopt the log-scaled mean value for the $\\alpha=2$ case, $k=9.5\\times10^{-4}$~Myr$^{-1}$~pc$^2$~$M_{\\odot}^{-1}$.  With $k$ adopted, we now determine the gas column density where $t_0 = t_{SF}$, as this is the regime where significant gas depletion is expected in the data and thus where deviations from a power law star-gas correlation would be detectable.  For $\\alpha=2$, $t_0$ is the time at which $\\Sigma_{gas}(x,y,t_0) = 0.5~\\Sigma_{gas}(x,y,0)$.  Therefore:  \\begin{equation} t_0 = \\frac{1}{2k} \\Sigma_{gas}(x,y,t_0)^{-1} \\end{equation}  \\noindent Setting $t_0=t_{SF}$ and thus this equation equal to Eqn.~\\ref{tsfdef} and solving numerically, we find $\\Sigma_{gas}(x,y,t_{SF}=t_0) = 300 M_{\\odot}$~pc$^{-2}$.  Thus gas depletion only becomes noticeably relevant at gas column densities on the high end of the observed range.  In summary, the model closely follows the observed power law correlation between star formation rate surface density and gas column density.  Only at the highest observed column densities does the model predict significant deviations from a power law: these may be undetected due to the scatter in the observed correlation.  At this point, it is useful to visualize the time evolution of the simple gas depletion model and compare it with our observations.  In Fig.~\\ref{alpha2}, we plot the realizations of Equation~\\ref{fullsoln} for several values of $t$ in the range of 200 years to 6.3~Myr and a range of $\\Sigma_{gas}(t=0)$ from 20 to 2000~$M_{\\odot}$~pc$^{-2}$, assuming $\\alpha$ = 2, $c$ = 0.5, and $k = 9.5 \\times 10^{-4}$~Myr$^{-1}$~pc$^2~M_{\\odot}^{-1}$.  These realizations are isochrones showing the locus in the $\\Sigma_*$ vs. $\\Sigma_{gas}$ diagram for an ensemble of co-eval molecular cloud parcels with different initial values of $\\Sigma_{gas}$.  We also plot orthogonal tracks that trace the evolution in $\\Sigma_*$ vs. $\\Sigma_{gas}$ for parcels of a molecular cloud as the gas of a given $\\Sigma_{gas}(t=0)$ is converted into stars.  The polygonal regions that demarcate the observed data limits for MonR2 (green) and Ophiuchus (blue) are overplotted for comparison.  At early $t$, the model follows a power law of index $\\sim$2.  As we approach typical star formation region ages of $\\sim$1~Myr, the highest gas density evolutionary tracks curve toward decreasing gas densities, resulting in local steepening of the model isochrone. This is a direct result of the varying relative rate of gas depletion at differing column densities.  For the lowest initial gas densities, the evolutionary tracks follow an almost vertical trajectory as stars are formed at very low efficiency.  In contrast, the higher initial gas densities form stars very efficiently and thus the gas is quickly depleted relative to the star formation time, $t_{SF}$.  We can easily rewrite Eqn.~\\ref{sigrat} to describe the star formation efficiency at $t=t_{SF}$ as:  \\begin{equation} \\frac{\\Sigma_*}{\\Sigma_*+\\Sigma_{gas}}(x,y,t_{SF}) = \\frac{\\Sigma_{gas}(x,y,t_{SF})}{\\Sigma_{gas}(x,y,t_{SF})+1/(ckt_{SF} (1+\\frac{t_{SF}}{t_0}))} \\end{equation}  \\noindent At the $\\Sigma_{gas}$ fiducial points of 20 and 300~$M_{\\odot}$~pc$^{-2}$ adopted above, $\\frac{\\Sigma_*}{\\Sigma_*+\\Sigma_{gas}}(x,y,t_{SF})~=~2.3\\%$ and 26\\% respectively.  These are similar to reported values for ``distributed'' and ``clustered'' star formation \\citep[e.g.][]{carp00,alle07,evan09}.  In this analysis, as is claimed by \\citet{bres10}, we see a smooth variation from low to high column density star formation.  We add to that analysis that the most likely progenitor of the smooth variation in stellar densities is the smooth variation of gas column densities in star forming clouds.  In Fig.~\\ref{alpha}, we investigate different values of $\\alpha$, as well as the possibility of a threshold in density for star formation, and compare the results to the fits to the MonR2 and Ophiuchus clouds.  We show Equation~\\ref{fullsoln} models for $\\alpha =1$ (top left panel) and $\\alpha = 1.4$ (top right panel), with no threshold imposed, and with appropriate values of $k$ ($\\alpha$-dependent) to approximately fit the MonR2 and Ophiuchus data (green and blue polygons, respectively): neither provides a good fit to the data, unlike that in Fig~\\ref{alpha2}.  We also implement a threshold (bottom two panels) such that the star formation ceases when $\\Sigma_{gas}(x,y,t)$ is less than a value of 100~$M_{\\odot}$~pc$^{-2}$.  A less extreme threshold has been reported in a different analysis of the {\\it Spitzer} c2d and Gould Belt Legacy survey data, where protostars alone are used to compute star formation surface densities \\citep{heid10}.  Our analysis is not consistent with a star formation threshold at any gas column density in the range that we have surveyed.   \\subsection{A Simple Model of Jeans Fragmentation\\label{model}}  For MonR2 and Ophiuchus, the power law indices in the relation between stellar and gas surface density are close to 2.  This section shows that such an exponent is expected from a simple model of Jeans fragmentation in an isothermal self-gravitating layer.  If one initial Jeans mass makes one star, then the number surface density of stars in a layer equals the number surface density of Jeans masses,  \\begin{equation} N_* =  N_J \\label{jeans1} \\end{equation}  \\noindent The Jeans mass in a self-gravitating layer is  \\begin{equation} M_J = \\frac{B \\sigma^4}{\\Sigma_{gas}~G^2} \\label{jeans2} \\end{equation}  \\noindent where $B$ is a constant \\citep[4.67 for an isothermal, self-gravitating layer of gas;][]{lars85}, $\\sigma$ is the gas velocity dispersion, $\\Sigma_{gas}$ is the gas mass surface density, and $G$ is the gravitational constant.  \\noindent Then the number surface density of Jeans masses can also be written as  \\begin{equation} N_J  = \\frac{\\Sigma_{gas}}{M_J} \\end{equation}  \\noindent or substituting from \\ref{jeans2},  \\begin{equation} N_J = \\frac{1}{B} (\\frac{G \\Sigma_{gas}}{\\sigma^2})^2 \\end{equation}  \\noindent and substituting from \\ref{jeans1}, multiplying both sides of the equation by a mean mass per star of 0.5~$M_{\\odot}$\\footnote{Note that the eventual mass of a given star is explicitly assumed to be uncorrelated with the local Jeans mass.  This is consistent with some characterizations of the pre-stellar core to stellar mass conversion process that include potential accretion of gas from the larger-scale environmental distribution \\citep[e.g.][]{myer08,myer09b}.}, and converting to mass surface density units of $M_{\\odot}$~pc$^{-2}$, with an assumed $\\sigma = 0.2~$km~s$^{-1}$, the sound speed at 10~K:   \\begin{equation} \\Sigma_* = 1.34 \\times 10^{-3} \\Sigma_{gas}^2 \\end{equation}  Thus the power law of index 2 is predicted by simple thermal fragmentation of an isothermal self-gravitating layer.  Furthermore, the value of $1.34 \\times 10^{-3}$ is in the middle (in log-space) of our observed range of $\\Sigma_*/\\Sigma_{gas}^2$ from section~\\ref{corr}, and represents a mean value for $ckt(1+t_{SF}/t_0)$ in the evolutionary model above.  We note that there is no explicit time dependence in this model, thus although it gives the correct star-gas correlation, it does not predict a rate of star formation.  More generally, one must consider cloud layers that are non-uniform, spanning a wide range of column densities in order to replicate the trends we report here, which are found within each star-forming cloud individually.  An example of a horizontally nonuniform planar system is the density-modulated isothermal layer \\citep{schm67}, or of a planar system of isothermal cylinders \\citep{ostr64} can be expected to give a similar relation between stellar and gas surface densities.  Some of the clouds studied in this paper may have such a flattened geometry; their extinction maps indicate that ``hub-filament'' structure is expected if a clumpy medium is compressed into a modulated self-gravitating layer \\citep{myer09}.  If this picture of fragmentation in a layer applies to some of the clouds considered here, the results of Section~\\ref{tommodel} indicate that evolutionary effects probably do not change the relation between star and gas surface density very much, over the likely star-forming life of the cloud.   \\subsection{Empirical Comparison to Cloud-Scale Values, the Kennicutt-Schmidt Law, and Recent Parsec-scale Work}  As already noted, \\citet{evan09} compared the {\\it overall} star formation surface density vs. gas column density for each of the clouds in the c2d survey, as well as an average of all of them, for comparison with the Kennicutt-Schmidt law for galaxies.  They find that the star formation rate density predicted by the law falls a factor of 20 below their observed data. In Figure~\\ref{neal1}, we have overlaid the correlation reported here for even smaller scales (0.3 to 2.0~pc) on our reproduction of the \\citet{evan09} figure.  Our correlation overlays the \\citet{evan09} cloud-averaged data well at intermediate surface densities, and the mean value for all clouds falls quite close to our correlation.  The difference is only clear where our trend extends to higher column densities.  Because we are able to probe a wider range in gas column density by taking measurements at smaller size scales, it is not surprising that such a discrepancy could be uncovered.  Regardless, our analysis agrees with that of \\citet{evan09}; studies of star formation rate density versus gas density for nearby clouds show much greater star formation rate for the local gas density than extragalactic correlations would suggest. As noted by \\citet{evan09} and others, the measurements used to derive most star formation law characterizations vary widely in the size scales and time scales probed, and these variations almost certainly play a role in the apparent discrepancies.  Indeed, recent work to characterize the role of measurement size scale in characterizing the correlation between star formation rate surface densities and gas surface densities in nearby galaxies has revealed considerable change in the power law index with the measurement size scale used \\citep{liu11}.  Other recent efforts have examined the star formation rate and gas surface density correlation in nearby star-formation regions via similar sorts of data (ie. YSO star counts and near-IR extinction maps) and some overlap in the specific clouds considered. \\citet{lada10} estimated the star formation rate vs. cloud mass for a sample of 11 nearby clouds.  Within this cloud sample, they noted that there was a wide dispersion in the estimated star formation rate per cloud mass.  They then demonstrated that the observed dispersion  within their cloud sample was minimized if they only considered the cloud mass at column densities above $A_V = 7$ (or 116~$M_{\\odot}$~pc$^{-2}$).  They proposed that this corresponded to a density threshold for star formation; for densities above this threshold the star formation rate varied linearly with gas mass, below the threshold the star formation rate was negligible.  A similar threshold ($129 \\pm 14~M_{\\odot}$~pc$^{-2}$) was invoked by the sub-cloud-scale analysis by \\citet{heid10}, a follow-up to the cloud-scale examination by \\citet{evan09} mentioned above.  \\citet{heid10} used the numbers of {\\it Spitzer}-identified Class~I and so-called ``flat spectrum'' YSOs (ie. they excluded the longer-lived Class~II YSOs) within various extinction contour intervals to probe the correlation in 20 nearby molecular clouds. Since their {\\it Spitzer} data did not include clouds forming massive stars, they did not have a direct measurement of YSO densities in high mass star forming regions. Thus they added star formation rates derived from {\\it IRAS} far-IR luminosities and mean gas column densities derived from HCN line emission for a set of more distant and presumably more active star-forming regions.  We have overlaid their final bent power law best fit to their combined data in Fig.~\\ref{neal1} (red solid lines), showing that it underestimates the star formation rates we report here at both extremely high and low gas column densities.  At high column densities, the shallow portion of the bent power law of \\citet{heid10} is driven by the inclusion of the {\\it IRAS}-HCN-derived data in the fit; their {\\it Spitzer}-derived data seem relatively consistent with our analysis (examined in more detail below).  That work includes a discussion of the likelihood that far-IR-derived star formation rates may be underestimated for individual star-forming regions by up to an order of magnitude.  We find plenty of evidence for that to be the case in the nearby embedded clusters, where we can do direct comparisons between the {\\it IRAS} far-IR-derived star formation rate and that derived from the YSO counts.  For example, using the known far-IR luminosity \\citep[26,000~$L_{\\odot}$;][]{ridg03} and the YSO count for the cluster core of the MonR2 main cloud core \\citep[132 YSOs, 0.73~pc effective radius;][]{gute09}, we obtain star formation rate surface densities of 3 and 20~$M_{\\odot}$~yr$^{-1}$~kpc$^{-2}$, respectively.  A similar discrepancy can be found in the Orion Nebula Cluster.  Its far-IR luminosity is estimated at $10^5~L_{\\odot}$ \\citep{geza98}, yet its central parsec contains of order $10^3$ YSOs \\citep[e.g.][]{feig05}, yielding SFR surface density estimates of 20 and $300~M_{\\odot}$~yr$^{-1}$~kpc$^{-2}$, respectively.  In summary, the common conversion of far-IR flux to star formation rate appears to underestimate the true star-formation rate at parsec scales in nearby regions.    We are not concerned with the apparent discrepancy with our result at low gas column densities. The protostellar number counts \\citep{heid10} at these gas column densities are very low: a star surface density of $10^{-1.7}~M_\\odot$~kpc$^{-2}$~yr$^{-1}$ corresponds to one $0.5~M_\\odot$ protostar per 0.5~Myr (the estimated lifetime of the evolutionary phase) in the area of an entire typical molecular cloud ($\\sim10$~pc size scale is typical of nearby clouds). We used the fitting technique described in Section~\\ref{corr} to fit their {\\it Spitzer}-derived data, excluding the upper limits.  We find that the {\\it Spitzer}-only correlation to be a power law of slope $3.7 \\pm 1.6$ (dashed red line in Fig.~\\ref{neal1}), a steeper power law than our reported trend (though within $1\\sigma$) and lacking in an apparent gas surface density threshold for star formation.  Differences at this level may be explained in part by differences in source classification method and uncertainty in the relative durations of the pertinent YSO evolutionary stages, but detailed consideration of such effects is beyond the scope of this paper."
    },
    "1107/1107.2977_arXiv.txt": {
        "abstract": "We analyze 26 archival Kepler transits of the exo-Neptune HAT-P-11b, supplemented by ground-based transits observed in the blue (B-band) and near-IR (J-band).  Both the planet and host star are smaller than previously believed; our analysis yields $R_p=4.31R_{\\oplus}\\pm0.06R_{\\oplus}$, and $R_s = 0.683R_{\\odot}\\pm 0.009 R_{\\odot}$, both about $3\\sigma$ smaller than the discovery values.  Our ground-based transit data at wavelengths bracketing the Kepler bandpass serve to check the wavelength dependence of stellar limb darkening, and the J-band transit provides a precise and independent constraint on the transit duration.  Both the limb darkening and transit duration from our ground-based data are consistent with the new Kepler values for the system parameters. Our smaller radius for the planet implies that its gaseous envelope can be less extensive than previously believed, being very similar to the H-He envelope of GJ\\,436b and Kepler-4b. HAT-P-11 is an active star, and signatures of star spot crossings are ubiquitous in the Kepler transit data.  We develop and apply a methodology to correct the planetary radius for the presence of both crossed and uncrossed star spots. Star spot crossings are concentrated at phases -0.002 and +0.006. This is consistent with inferences from Rossiter-McLaughlin measurements that the planet transits nearly perpendicular to the stellar equator.  We identify the dominant phases of star spot crossings with active latitudes on the star, and we infer that the stellar rotational pole is inclined at about $12^{\\circ}\\pm5^{\\circ}$ to the plane of the sky.  We point out that precise transit measurements over long durations could in principle allow us to construct a stellar Butterfly diagram, to probe the cyclic evolution of magnetic activity on this active K-dwarf star. ",
        "introduction": "The exo-Neptune HAT-P-11b (\\citealp{bakos}, hereafter B10) is prominent among extrasolar planets smaller than Saturn.  HAT-P-11b transits a bright star (V=9.59) that lies in the Kepler field \\citep{borucki}. Based on its position in a mass-radius diagram (e.g., Figure~14 of B10), HAT-P-11b is likely to have a massive atmosphere.  Moreover, B10 found good mass and radius agreement with metal-rich models for the planet \\citep{baraffe}, and the B10 spectroscopic analysis of the host star indicated that it was metal-rich.  The planet's atmosphere is therefore likely to exhibit a significant molecular absorption spectrum during transit and/or eclipse.  It is a tempting target for future spectroscopic characterization, for example using precise ground-based spectrophotometry (e.g., \\citealp{bean}) in combination with HST \\citep{pont09}, and/or Warm Spitzer \\citep{desert}. Prior to such efforts, it is important to improve our current knowledge of the system parameters and optical planetary radius by examining the Kepler data.  Based on photometric variations of the star, B10 concluded that star spots were common on the stellar photosphere. Rossiter-McLaughlin observations of the system \\citep{winn, hirano} indicate that the planet's orbital angular momentum vector is nearly perpendicular to the orbital angular momentum vector of the star.  \\citet{winn} predicted that this mis-alignment would produce a characteristic signature in the spot-crossing patterns seen during transit, and this should be quite evident in the Kepler data.  In this paper, we report an analysis of Q0-Q2 archival Kepler data for transits of HAT-P-11b, supplemented with new ground-based transit data at wavelengths bracketing the Kepler bandpass. The potential benefits of ground-based transit photometry as a complement to Kepler have been emphasized by \\citet{colon}. ",
        "conclusions": "Our MCMC fits produce excellent agreement with the Kepler data (Figure~6). The retrieved parameters agree closely whether we solve for limb darkening, or fix the coefficients at their Kurucz model atmosphere values.  Moreover, the retrieved linear coefficient in the former case ($u = 0.626 \\pm0.014$, Table~1) is in excellent agreement with the model atmosphere value ($0.6179$).  We conclude that the HAT-P-11 system parameters derived from the Kepler MCMC fits are robust, and not model-dependent via limb darkening. \\citet{knutson} reached a similar conclusion in their analysis of HST observations of HD\\,209458b.  The $\\chi^2$ value for the best fit solution shown on Figure~6 is 417 for the in-transit points, for 112 degrees of freedom.  Thus, the error bars based purely on Kepler photometric precision must be increased by a factor of 2 to account for the imperfect precision of star spot removal.  That factor was applied in the MCMC fits, as noted in Sec.~5.2.  Both the J-band and B-band transits imply linear limb darkening coefficients ($u$ in Table~1) that are in reasonable agreement with model atmosphere values, but they do hint that limb darkening for the real star could vary more strongly with wavelength than the model atmosphere predicts.  The MCMC posterior distribution for $u$ in the B-band (not illustrated) peaks at unity, and values exceeding unity are unphysical. The 1-sided Gaussian distribution has $\\sigma = 0.080$, so our result (Table~1, $u=1.000 \\pm0.080$) falls {\\it above} the model atmosphere value ($0.862$) by $1.8\\sigma$.  The J-band result ($u=0.086\\pm0.065$) similarly falls {\\it below} the model atmosphere value ($0.244$) by $2.4\\sigma$.  If discrepancies in model atmosphere predictions for $u$ do vary with wavelength in this fashion (stronger at short-$\\lambda$, weaker at long-$\\lambda$), we would not necessarily expect a significant effect in the Kepler band, because it lies intermediate in wavelength between our B- and J-band data. Moreover, any such systematic variation should be confirmed using transits of larger planets, exhibiting deeper transits, where greater precision in derived limb darkening can be achieved. We note that there is observational precedent for limb darkening at short wavelengths to be stronger than model atmosphere predictions \\citep{tingley}. We plan additional simultaneous B- and J-band observations, of giant planet transits.  Our results for the J-band transit are inconsistent with the Kepler results as regards the planetary radius.  In J-band, we find that $R_p/R_s$ is 6\\% larger than the Kepler solution ($0.0627$ {\\it vs.} $0.0589$), and the difference is more than 5 times the precision of the J-band measurement. We regard the Kepler result as definitive, so we consider how this discrepancy can be explained.  One potential explanation of a discrepant radius in J-band is that it reflects a true variation of the planetary radius with wavelength, due to atmospheric opacity.  However, the difference here seems implausibly large, and we prefer a more mundane explanation. The transit of HAT-P-11b is relatively long in duration (2.3 hours), and shallow (0.004).  Ground-based infrared photometry can be subject to baseline fluctuations caused by telluric water vapor absorption at the edges of the JHK bandpasses.  The longer the duration of a transit event, the more sensitive it is to baseline effects, because the adopted baseline has to span a longer interval.  Also, a given baseline error will have a greater relative effect for shallow transits.  Hence, HAT-P-11 is particularly prone to baseline errors, and we regard the $R_p/R_s$ value from our J-band MCMC fits as unreliable at the level of accuracy needed for meaningful comparison with Kepler results. We note that the discrepancy would be even larger without the baseline correction that we applied in Sec.~2.2; evidently our correction underestimates the telluric effects.  Although $R_p/R_s$ from our J-band data is questionable compared to the Kepler value, we believe that $a/R_s$ is reliable, and a useful complement to the Kepler results.  As noted in Sec.~5.4, the J-band transit is sensitive to the total chord length.  Already we have seen (Figure~1), that the duration of the J-band transit predicted from the B10 discovery results (blue curve on Figure~1) does not fit the transit duration seen in the J-band data.  Unlike the situation with transit depth, the transit duration observed in J-band is insensitive to the telluric atmosphere.  At this wavelength, the ingress and egress are sharp and well defined, and rapid changes of this type are much less likely to be caused by the telluric atmosphere.  Fixing the orbital inclination at the Kepler value, our J-band fits give $a/R_s =16.454\\pm0.131$, in excellent agreement with the Kepler value of $16.549\\pm0.230$.  Transit determination of $a/R_s$ can be used to derive the stellar density \\citep{seager}, for comparison to asteroseismology results.  However, the Kepler asteroseismology results for HAT-P-11 were regarded as preliminary by \\citet{cd}, so the comparison is premature.  We adopt our Kepler-1 solution (see Table~1) for the system parameters, and note that the stellar mass ($0.809M_{\\odot}$)is well-determined from the Hipparcos parallax and the isochrone fits performed by B10.  This yields $R_s = 0.683R_{\\odot}\\pm 0.009 R_{\\odot}$, a $3.3\\sigma$ revision from the B10 value ($0.752 R_{\\odot}\\pm0.021R_{\\odot}$).  Using our $R_p/R_s = 0.05892\\pm0.00027$, yields $R_p=4.39R_{\\oplus}\\pm0.06R_{\\oplus}$, a $2.1\\sigma$ revision from B10 ($R=4.73\\pm0.16R_{\\oplus}$).  However, the presence of unseen star spots will cause the planet to appear larger in the transit solutions by 1.76\\%, as discussed in Sec.~7.  We thus correct our radius for the planet downward to $R_p=4.31R_{\\oplus}\\pm0.06R_{\\oplus}$. Hence we find that both the planet and star are smaller than the B10 discovery values, at about the $3\\sigma$ level.  We have been able to improve the system parameters of HAT-P-11 over the very thorough analysis by B10, in large part due to the numerous precise Kepler transits.  These data have allowed us to remove the effect of star spot crossings, resulting in a deeper transit. Although the deeper transit would tend to produce a larger planet, we also find a smaller star and the decrease in stellar size, combined with our correction for uncrossed star spots, results in a smaller planet compared to B10.  Also in comparison to B10, we find a longer transit duration, and a smaller impact parameter since our orbital inclination is closer to 90 degrees.  Our revised radius for HAT-P-11b has noteworthy implications for its internal structure.  A mass-radius diagram for planets similar to HAT-P-11b is illustrated by \\citet{lissauer} (their Fig.~5). The H-He envelope mass implied by the previous radius of HAT-P-11b is close to 20\\%; our new radius reduces the required envelope mass to about 14\\%, very close to the envelopes implied for GJ\\,436b and Kepler-4b.  The radius of HAT-P-11b is now known with sufficient precision to contemplate useful comparisons with other space-borne transit observations, such as HST and Spitzer could acquire.  That comparison could potentially reveal an atmospheric signature of HAT-P-11b. Moreover, we point out that long-term monitoring of the system by Kepler, potentially supported by precise ground-based photometry after Kepler has ended, could reveal fundamental information on the magnetic cycle of the star.  Because the planet transits nearly perpendicular to stellar active latitudes, monitoring of the number and orbital phase of star spot crossings could allow us to construct a stellar `Butterfly diagram' depicting the cyclic evolution of magnetic eruptions at the stellar photosphere of this active K-dwarf star.  While this paper was under review, two additional analyses of HAT-P-11b based on the Kepler data were announced. \\citet{southworth11} obtains system parameters that agree with our Table~1.  He does not explicitly correct for the effect of star spots, instead treating the spot perturbations as correlated errors.  An analysis by \\citet{sanchis-ojeda} also finds similar system parameters as per our Table~1.  These authors correct for crossed spots `by hand', but do not correct the planetary radius for the effect of uncrossed spots. Interestingly, \\citet{sanchis-ojeda} also discuss a second possible orientation for the system in which the star is viewed nearly pole-on. In that orientation, the two preferred phases of star spot crossings could be caused by one active band of spots encircling the stellar pole.  While we cannot rigorously exclude that geometry, we note that our correction for uncrossed spots (Sec.~7) - based on viewing the star as per the Figure~8 geometry - predicts a peak-to-peak variation in the stellar rotational light curve of $\\leq 1.76\\%$, in good agreement with the variability seen in the Kepler data ($\\sim 1.5\\%$). That agreement would have to be accidental if the star is being viewed pole-on, so we believe that our Figure~8 geometry is correct, and the stellar pole is nearly in the plane of the sky."
    },
    "1107/1107.0005_arXiv.txt": {
        "abstract": "It is suspected that the ultraviolet (UV) upturn phenomenon in elliptical galaxies and extended horizontal-branch stars in globular clusters have a common origin. An extremely high abundance of helium ($Y \\sim 0.4$) allows for a working hypothesis, but its origin is unclear. Peng \\& Nagai (2009) proposed that primordial helium sedimentation in dark haloes over cosmic timescales may lead to extreme helium abundances in galaxy cluster centers. In this scenario UV upturn should be restricted to brightest cluster galaxies (BCGs) only. This is a clear and testable prediction. We present tests of this hypothesis using galaxy clusters from Yoon et al. (2008) that were detected by both the Sloan Digital Sky Survey and the Galaxy Evolution Explorer Medium Imaging Survey. Using a new UV classification scheme based on far-UV, near-UV, and optical photometry we found only 5\\% of cluster elliptical galaxies show a UV upturn, while 27\\% and 68\\% are classified as ``recent star-formation'' and ``UV-weak'' ellipticals, respectively. The data reveal a modest positive dependence of the UV upturn fraction on galaxy velocity dispersion, which is in agreement with the earlier findings of Burstein et al. (1988) and possibly with the helium sedimentation theory. However, we do not see any dependency on rank or luminosity of galaxies. Besides, BCGs do not show any marked difference in UV upturn fraction or strength, which is inconsistent with the prediction. We conclude that the aforementioned helium sedimentation theory and its inferred environmental effects are not supported by the available data. ",
        "introduction": "The ultraviolet (UV) upturn phenomenon in elliptical galaxies has garnered much attention since its discovery (Code \\& Welch 1979; Burstein et al. 1988; see O'Connell 1999 for a review). When the phenomenon was first observed, it was considered a mystery because elliptical galaxies were thought to be almost exclusively composed of old stars. Such stars were believed to be poor UV sources. The situation has significantly changed in recent years. Many elliptical galaxies still seem to form stars at low levels (Yi et al. 2005), and even old stars can be effective UV producers (Greggio and Renzini 1990; Horch et al. 1992; Bressan et al. 1994; Dorman, O'Connell, and Rood 1995; Brown et al. 2000; Buzzoni 2007). Stellar population synthesis models based on the old horizontal-branch star assumption appear to provide the most compelling explanation for UV upturn. However, it is still unclear exactly how some, but not all elliptical galaxies develop such a large number of UV bright stars and exhibit a UV upturn.  Recent research focused on extreme helium enrichment has provided an interesting possibility. In studies on globular clusters, some massive globular clusters, most notably $\\omega$~Cen, have been found to exhibit multiple main sequences (Bedin et al. 2004) and interestingly, the bluest main sequence seems to be the most metal-rich (Piotto et al. 2007). This is possible if the most metal-rich stars are extremely helium-rich as well (Norris 2004). Helium is extremely difficult to measure directly and so these claims are dependent upon {\\em indirect} deduction based on stellar colors and metal line strengths. An inferred helium abundance of $Y \\gtrsim 0.4$ apparently explains the presence of an extended horizontal branch of hot temperatures (Lee et al. 2005). Such UV-bright globular clusters have been found in the elliptical galaxy, Cen A, as well (Sohn et al. 2006; Kaviraj et al. 2007b). This helium abundance value corresponds to $\\Delta Y \\gtrsim 0.16$ when assuming $Y \\approx 0.24$ for non-helium-enhanced stars and poses a serious theoretical challenge with respect to helium-to-metal enrichment scenarios (Choi and Yi 2007; 2008). Internal stellar evolution theory alone does not seem to present a clear solution. Intricate dynamic calculations (D'Ercole et al. 2008; Decressin, Baumgardt, \\& Kroupa 2008) have been introduced to provide a working hypothesis.   An alternative solution has been presented by Chuzhoy and Loeb (2004) and Peng and Nagai (2009). Their theoretical models based on primordial helium sedimentation suggest that the environment has a profound effect on the UV spectral evolution of elliptical galaxies. In this work, we focus on the Peng and Nagai study because it has a more dramatic effect. The sedimentation effect is suspected to be larger for more massive galaxies and most extreme for galaxies at the center of galaxy clusters. According to Peng and Nagai (2009), this effect can raise the helium abundance {\\em of the gas} in individual galaxies near their center by as much as $\\Delta Y_{\\rm gal} \\sim 0.1$ over a Hubble time, which by itself does not seem large enough to explain the extreme horizontal branch and the UV upturn.  Sedimentation is thought to be far more dramatic up to $\\Delta Y_{\\rm clu} \\sim 0.4$ towards the galaxy cluster center. This effect is expected to be larger in more massive clusters. According to the theory, enrichment takes roughly 5 Gyr for $\\Delta Y_{\\rm clu} \\sim 0.2$ in massive clusters. If a substantial fraction (e.g., $\\sim 30$\\%\\footnote{Such a large fraction (30\\%) of intermediate-age stars in BCGs hypothesized here is generally not supported.} as found in $\\omega$~Cen) of stars in the brightest cluster galaxies (BCGs) formed from helium-enriched gas roughly 5 Gyr after the initial starburst in the galaxy, it might reach $\\Delta Y_{\\rm BCG} = \\Delta Y_{\\rm gal} + \\Delta Y_{\\rm clu} \\gtrsim 0.2$  and explain the UV upturn.  In addition, the sedimentation effect is supposed to be larger for a more massive cluster halo. This theory allows for a clear and testable prediction: UV upturn should mostly be restricted to BCGs that had substantial residual star formation and is more pronounced in more massive cluster haloes.  It is unclear whether the UV strengths of globular clusters have any causal connection with the UV upturn phenomenon in elliptical galaxies, as they occupy widely separated parameter spaces. However, both communities raising the same issue of extreme helium abundance seems to call for an investigation. We hereby perform a test on the above prediction using Sloan Digital Sky Survey (SDSS) optical and Galaxy Evolution Explorer (GALEX) UV photometric data. ",
        "conclusions": "In the Peng \\& Nagai scenario, the UV upturn is the result of three essential elements: (1) the assumption that extreme-helium stars evolve into UV-bright phases, (2) helium sedimentation, and (3) prolonged star formation. Discussing the first condition is beyond the scope of this paper anyway. Considering the amount of helium sedimentation alone, BCGs become the best candidate to show a UV upturn. The last element can be examined from the results of other studies.  Our test is based on a sample that is an order of magnitude larger than what was previously available. We also used a new {\\em empirical} classification scheme to identify UV upturn, RSF, and UV-weak galaxies.  In cluster environments, roughly 5\\% of elliptical galaxies are classified as UV upturn galaxies. Brighter galaxies appear to show a UV upturn more frequently, in agreement with an earlier report by Burstein et al. (1988), but the mass (in rank and brightness) dependence is weak at least for the narrow mass baseline of our sample. BCGs show a UV upturn slightly more frequently than fainter galaxies do, but this does not seem significant in any way as predicted by the Peng \\& Nagai scenario, especially when considering the underlying mass dependence.   If helium sedimentation did indeed occur over the Hubble timescale in all large haloes as predicted, this fraction (5\\%) appears to be too small. Such findings may mean that the majority of cluster elliptical galaxies,  especially BCGs, did not have an extended star formation history hence did not benefit from helium sedimentation. This is compatible with the general understanding on cluster elliptical galaxy evolution (e.g., Bower, Lucy, \\& Ellis 1992). It is then suggested that UV upturn galaxies are rare ellipticals with an extended star formation history. This is however in contrast to empirical findings. Burstein et al. (1988) found that UV stronger galaxies tend to be more abundant in magnesium, a representative alpha element. Such a finding was later confirmed by Bureau et al. (2011). As alpha enhancement is generally considered to be the result of a short starburst on the order 1~Gyr or less (e.g., Worthey, Faber, \\& Gonzalez 1992), such an enhancement seems to contradict the requirement of this particular hypothesis.   If helium sedimentation indeed took place as theory suggests, there might have been reasons for that BCGs do not appear special in the UV. First, some non-BCG satellite (but massive) ellipticals may have been BCGs in the previous cluster halo before halo mergers. For example, one can assume that more than a couple of cluster halo mergers occurred in the last few billion years to make one present-day cluster (which is probable). Assume that BCGs in the merging clusters are still orbiting as satellites in the merger remnant cluster. The lifetimes of the satellites that were formerly BCGs depend on the details of the merger conditions such as halo mass ratios and orbital parameters (Boylan-Kolchin, Ma \\& Quataert 2008; Kimm, Yi \\& Khochfar 2011), but they can be on the order of a few billion years. If so, the difference between the present-day BCGs and massive satellite ellipticals may not be as large as the closed-box assumption taken in our investigation. In this regard, one may consider all the bright satellite ellipticals in a cluster ``former BCGs''. Such ellipticals may have then benefitted from the helium sedimentation from their former halo environments and so exhibit a modest UV upturn. However, the Burstein et al. relation between magnesium abundance and UV strength still contradicts such a notion. Besides, the ``former BCG'' argument seems a bit too contrived.   We stress that our study takes a purely empirical approach. But this raises an issue regarding the time evolution of UV upturn. The mean redshift of our sample is roughly 0.07, which corresponds to a lookback-time of 0.9\\,Gyr from the local universe. Some theories of UV upturn (e.g., Yi et al. 1999; Buzzoni 2007) predict a fast evolution of UV strength with time, while different predictions are available too (e.g. Han et al. 2007). The fast-evolution models suggest that the mean UV strength may have been slightly lower at $z=0.07$ than today. We admit that our analysis could be subject to such details. But we believe that the difference in lookback-time in this case is not large enough to cause a major threat to the interpretation. Besides, while the redshift evolution might be real, we still do not see any environmental effect in our sample (at $z=0.07$) in the manner that has been predicted by the helium sedimentation theory.  We conclude that the our data are not compatible with the current version of  the helium sedimentation theory. Environmental effects on the UV strength of elliptical galaxies, if any, seem elusive."
    },
    "1107/1107.0719_arXiv.txt": {
        "abstract": "{ The Compact Symmetric Objects (CSOs) are small ($<1$\\,kiloparsec) and powerful extragalactic radio sources showing emission on both sides of an active galactic nucleus and no signs of strong relativistic beaming. They may be young radio sources, progenitors of large FR~II radio galaxies. } { We aim to study the statistical properties of CSOs by constructing and investigating a new large sample of CSO candidates on the basis of dual-frequency, parsec-scale morphology. } { For the candidate selection we utilized VLBI data for 4170 extragalactic objects obtained simultaneously at 2.3 and 8.6\\,GHz ($S$ and $X$ band) within the VLBA Calibrator Survey 1-6 and the Research and Development -- VLBA projects. Properties of their broad-band radio spectra were characterized by using RATAN-600 observations. Numerical modeling was applied in an attempt to explain the observed effects. } { A sample of 64 candidate CSOs is identified. The median two-point $S$--$X$~band spectral index of parsec-scale hotspots is found to be $-0.52$; with the median brightness temperature $\\sim 10^9$\\,K at $X$~band. Statistical analysis reveals a systematic difference between positions of brightest CSO components (associated with hotspots) measured in the $S$ and $X$~bands. The distance between these components is found to be on average $0.32 \\pm0.06$\\,mas greater at $8.6$\\,GHz than at $2.3$\\,GHz. } { This difference in distances cannot be explained by different resolutions at the $S$ and $X$~bands. It is a manifestation of spectral index gradients across CSO components, which may potentially provide important physical information about them. Despite our detailed numerical modeling of a CSO hotspot, the model was not able to reproduce the magnitude of the observed positional difference. A more detailed modeling may shed light on the origin of the effect. Multifrequency follow-up VLBI observations of the selected sample are needed to confirm and study newly identified CSOs from the list presented here. } ",
        "introduction": "Compact Symmetric Objects (CSOs) are small ($<1$\\,kiloparsec), powerful extragalactic radio sources that show emission on both sides of an active galactic nucleus \\citep{1994ApJ...432L..87W,1996ApJ...460..612R}. In contrast to the majority of compact radio sources, relativistic beaming effects are believed to be small in CSOs owing to their orientation close to the plane of the sky. The parsec-scale core that marks position of the central engine is often weak or not detected at all. Kinematic studies of CSOs reveal no superluminal motion and suggest source ages of a few hundred to a few thousand years \\citep{1998A&A...337...69O,2009AN....330..149P}. CSOs may be progenitors of large-scale Fanaroff-Riley type II radio galaxies \\citep{1995A&A...302..317F,1996ApJ...460..634R,2002ApJ...568..639P}. To draw solid conclusions about the properties of CSOs as a class, it is important to construct a large representative sample of them.  An interesting side effect of an investigation of a large CSO candidate sample is the possibility of finding a supermassive binary black hole pair \\citep{2004ApJ...602..123M,2006ApJ...646...49R,2009AN....330..206T} that can mimic a CSO. Such an object would have two compact radio components that share the typical characteristics of the core: flat spectrum, high brightness temperature $T_\\mathrm{b}$, and complex flux variability pattern. It has also been suggested that true CSOs are the likely hosts of supermassive black hole binaries \\citep{2010ApJ...713.1393W}. Finally, CSOs may be useful as calibrators for continuum observations of other radio sources thanks to their stable flux density and low fractional polarization \\citep{2003ApJ...597..157T}.  The largest homogeneously selected CSO sample to date is the COINS sample by \\cite{2000ApJ...534...90P}. An initial list of candidates for this sample was selected from \\cite{1988ApJ...328..114P} and Caltech--Jodrell Bank \\citep{1994ApJS...95..345T} VLBI surveys and  from the first VLBA Calibrator Survey \\citep{VCS1}. Dedicated multifrequency polarimetric VLBA observations of the candidates have been performed to distinguish between true CSOs and contaminating core-jet type sources. The sample was further extended to the northern and southern extremities of the VLBA Calibrator Survey by \\cite{2003ApJ...597..157T}. Another CSO sample based on the VLBA Imaging and Polarization Survey is being constructed by \\cite{2009AN....330..206T}.  In this paper we present a list of CSO candidates selected by using publicly available\\footnote{\\url{http://astrogeo.org/vlbi_images/}} $S$~band (central frequency 2.3\\,GHz) and $X$~band (central frequency 8.6\\,GHz) VLBI data collected in the course of the VLBA Calibrator Surveys (VCS) 1 to 6 by \\cite{VCS1}, \\cite{VCS2}, \\cite{VCS3,VCS4}, \\cite{VCS5}, \\cite{VCS6}, and the Research and Development VLBA program \\citep[RDV, e.g.,][]{1996ApJS..105..299F,1997ApJS..111...95F,2008evn..confE..86P,RDV2009}. We discuss the basic properties of the compiled sample of CSO candidates including an unexpected systematic difference in CSO component positions measured at the lower and higher frequencies and our attempt to reproduce this effect with detailed numerical modeling. ",
        "conclusions": "\\label{sec:discussion}  When the jet flow crosses the hotspot, it is deflected and starts to flow backwards towards the nucleus. This is known as the backflow. When radio-emitting particles cool in such a backflow, we expect the higher energy emission (caused by the faster cooling high-energy particles) to be concentrated closer to the hotspot, whereas the lower energy emission (from the slower cooling low energy particles) should be seen farther away towards the nucleus, while these particles are being dragged by the backflow. Thus, this should generate a spatial separation between the peaks of the emission at higher and lower frequencies such that the high-frequency peak is observed farther away from the nucleus than the lower frequency one. This is a possible cause of the observed bias.  We attempted to test this hypothesis by computing radio emission from simulated relativistic jets. The code and initial setup is similar to that of \\cite{2007MNRAS.382..526P}. We performed two simulations, one with a one-sided jet (M1) and a second one with a two-sided jet (M2), by changing the boundary conditions at the inlet. The radio emission is computed using the \\emph{SPEV} code \\citep{2009ApJ...696.1142M}.  The simulated jets are injected with an initial velocity $0.95$~c and the Mach number $1.66$. The density contrast between the jets and external medium at the injection point ($10$~pc from the nucleus) is $10^{-5}$, and their initial radius is $2$~pc. The jet composition is purely leptonic and the external medium is composed of the neutral hydrogen. We assume the profile of the external medium in a typical radio galaxy like 3C~31 \\citep{2007MNRAS.382..526P,2002ApJ...581..948H}.  The jets were evolved until they reach the distance of $110$~pc from the nucleus.  We used these as an input for the \\emph{SPEV} code, and produced synthetic radio images at $2.3$ and $8.6$\\,GHz, assuming the jet angle to the line of sight to be $70^\\circ$. Since the simulations do not include magnetic fields, for the purpose of computing the emission, we assumed a randomly oriented magnetic field in a fraction of equipartition with the thermal fluid, whereby the ratio of magnetic to internal energy density was fixed to $\\epsilon_B = 10^{-3}$. The relative properties of the emission at $2.3$ and $8.6$\\,GHz are expected to depend only weakly on the value of $\\epsilon_B$ because the radio frequencies are always below the cooling frequency for the particle distribution. The typical value of the magnetic field at the injection point is $\\approx 20 \\mu$G.  \\emph{SPEV} code uses Lagrangian particles as markers of the nonthermal particles responsible for the radio emission (details can be found in \\citealt*{2009ApJ...696.1142M}). In the current work we assume that these Lagrangian particles are inserted into the jet at the injection point and their spatial, temporal and spectral evolution is followed by taking synchrotron and adiabatic losses into account. We also have the possibility of accelerating additional particles in strong shock waves (e.g., Mach disk). We denote such models with a suffix ''s'' (M1s and M2s). For both types of particle injection (at the jet injection point or in the shocks), we assume that the total particle number density is proportional to the thermal fluid density and that the total particle energy density is proportional to the fluid internal energy density. We fix the power-law index of the nonthermal energy distribution at the time of injection to $2.25$. These assumptions allow us to determine the lower cutoff of the particle injection spectrum (see \\citealt*{2009ApJ...696.1142M} for more details). To determine the upper cutoff we equate the synchrotron cooling time with the electron acceleration time (see \\citealt*{MGA10} for a more detailed discussion of the upper cutoff).  The synthetic radio maps obtained from the simulations have much better resolution than typical VLBI observations. In order to be able to compare them directly we artificially degrade the quality of synthetic radio maps to that of the observations. To this end we compute degraded images assuming the source to be at two redshifts, $z = 0.04$ and $z = 1$. We note that the angular resolution of the original synthetic radio maps is $0.14$ and $0.014$ milliarcseconds at $z=0.04$ and $z=1$, respectively. We convolve the images with a Gaussian beam whose characteristic half-power size is $6$ and $1.8$ milliarcesonds at $2.3$ and $8.6$\\,GHz, respectively. In Fig.~\\ref{fig:synthetic} we show radio maps for model M1.  \\addtocounter{table}{1} % \\begin{table} { \\caption{Difference in positions of the peaks at $8.6$ and $2.3$\\,GHz.} \\label{table:1} \\begin{center} \\begin{tabular}{lcc} \\hline\\hline Model & $\\Delta_{0.04}$ & $\\Delta_{1}$\\\\ index & (mas)           & (mas)       \\\\ \\hline M1s   & $<0.14$ & $0.11$  \\\\ M1    & $<0.14$ & $0.11$  \\\\ M2s   & $<0.14$ & $0.11$  \\\\ M2    & $<0.14$ & $0.11$  \\\\ M2cs  & $-0.28$ & $0.055$ \\\\ M2c   & $-0.28$ & $0.055$ \\\\ \\hline \\end{tabular} \\end{center} } \\textbf{Note:} $\\Delta_{0.04}$ and $\\Delta_1$ are the differences for the source redshift of $0.04$ and $1$, respectively. Positive values indicate that the peak at the $8.6$\\,GHz is farther away from the nucleus than at $2.3$\\,GHz. M1s and M1 stand for the one-sided jet model with and without shock acceleration, respectively. Similarly, M2s and M2 stand for the jet of the double-sided jet model, while M2cs and M2c stand for the counter jet in that same model. When the difference is smaller than the angular resolution of the unconvolved maps, we give only upper bounds. \\end{table}  \\begin{figure*}[t] \\centering \\includegraphics[scale=0.33,trim=1cm 0cm 0cm 0cm]{M2s_far_new.eps} \\caption[]{ Synthetic radio maps at $2.3$\\,GHz (lower panel) and $8.6$\\,GHz (upper panel) of the simulated jet model M2s (see Fig.~2 in \\citealt*{2007MNRAS.382..526P} for the typical hydrodynamic structure of such simulated jets). Shades of gray show the normalized intensity of the unconvolved radio maps. Assuming that the source is located at $z=1$ contours show the degraded (convolved) image. Contour levels are 1, 2, 4, 8, 16, 32, and 64\\,\\% of the peak intensity. The bright feature in the center of each panel is the hot spot, and vertical lines denote the longitudinal position of the peak intensity in the convolved image. The difference in position of the peak at two frequencies is very small.} \\label{fig:synthetic} \\end{figure*}  In Table~\\ref{table:1} we summarize the results of our simulations. We see that the results do not confirm the hypothesis. Although models M1, M1s, M2, and M2s, as well as models M2c and M2cs for $z = 1$, show the expected trend, the differences are too small to be detected, which can also be seen in Fig.~\\ref{fig:synthetic}, where vertical lines denote the position of the intensity peak. No such trend is observed for models M2c and M2cs for $z = 0.04$.  The population of the nonthermal particles is dominated by the synchrotron cooling close to the injection site and by the adiabatic cooling further downstream (see e.g., \\citealt*{2009ApJ...696.1142M}). In the synchrotron-cooling region we expect that the higher energy particles lose their energy faster than the lower energy ones, whereas in the adiabatic-cooling region they all lose the energy at the same rate. Nonthermal particles inserted at the jet injection point reach the adiabatic-cooling region long before they reach the hotspot. Although there is substantial compression of the flow in the shock, if shock acceleration is not included, this compression alone is not enough to produce another synchrotron-cooling region. Therefore, the cooling rate of the high and low energy particles is the same, and no spatial separation between high and low frequency emission is expected. If we include shock acceleration, there is a synchrotron cooling zone behind the shock, but it is not large enough to reach into the backflow, because the shock does not accelerate particles to high enough energy to produce a larger synchrotron cooling dominated zone. Therefore, the result is qualitatively the same as when shock acceleration is not taken into account. This can be seen in the unconvolved radio maps (Fig.~\\ref{fig:synthetic}), where no noticeable difference between the positions of the peaks can be observed.  The systematic difference in hotspot positions obtained in our model (Table~\\ref{table:1}) results fro the fact that the convolution beam at $2.3$\\,GHz is larger than at $8.6$\\,GHz. The former includes more emission from jet and backflow regions closer to the nucleus (blending), and this generates a shift in the position of the apparent peak towards the nucleus. This results in the aforementioned apparent position difference in our models, which is an effect of model image convolution necessary to make it comparable to real observations. The same effect should be present in observations of real sources. However, the predicted magnitude of this blending effect is still smaller than the observed position difference. This is in accordance with minor difference between the modeling results obtained using full (different resolution) and restricted (matched resolution) $uv$-range as discussed in Section~\\ref{s:pos}.  In future work we intend to study two possibilities that could explain the observations. On the one hand, a much more efficient shock acceleration could produce enough high-energy particles to create a longer synchrotron cooling zone. On the other hand, the process of shear layer acceleration (e.g., \\citealt*{2002ApJ...578..763S,2008ApJ...681...84A}) could produce continuous injection of nonthermal particles into the backflow. Small velocity gradients present there would probably result in low acceleration efficiency, lower emission frequency, and ultimately, a shift in the position of the peak of emission towards the nucleus for those frequencies."
    },
    "1107/1107.1235_arXiv.txt": {
        "abstract": "I calculate the physical properties of 32 transiting extrasolar planet and brown-dwarf systems from existing photometric observations and measured spectroscopic parameters. The systems studied include fifteen observed by the \\corot\\ satellite, ten by \\kepler\\ and five by the {\\it Deep Impact} spacecraft. Inclusion of the objects studied in previous papers leads to a sample of \\reff{58} transiting systems with homogeneously measured properties. The \\kepler\\ data include observations from Quarter 2, and my analyses of several of the systems are the first to be based on short-cadence data from this satellite.  The light curves are modelled using the {\\sc jktebop} code, with attention paid to the treatment of limb darkening, contaminating light, orbital eccentricity, correlated noise, and numerical integration over long exposure times. The physical properties are derived from the light curve parameters, spectroscopic characteristics of the host star, and constraints from five sets of theoretical stellar model predictions. An alternative approach using a calibration from eclipsing binary star systems is explored and found to give comparable results \\reff{whilst imposing a much smaller computational burden}.  My results are in good agreement with published properties for most of the transiting systems, but discrepancies are identified for \\corot-5, \\corot-8, \\corot-13, Kepler-5 and Kepler-7. Many of the errorbars quoted in the literature are underestimated. Refined orbital ephemerides are given for \\corot-8 and for the \\kepler\\ planets. Asteroseismic constraints on the density of the host stars are in good agreement with the photometric equivalents for HD\\,17156 and TrES-2, but not for HAT-P-7 and HAT-P-11.  Complete error budgets are generated for each transiting system, allowing identification of the observations best-suited to improve measurements of their physical properties. Whilst most systems would benefit from further photometry and spectroscopy, HD\\,17156, HD\\,80606, HAT-P-7 and TrES-2 are now extremely well characterised. HAT-P-11 is an exceptional candidate for studying starspots. The orbital ephemerides of some transiting systems are becoming uncertain and they should be re-observed in the near future.  \\reff{The primary results from the current work and from previous papers in the series have been placed in an online catalogue, from where they can be obtained in a range of formats for reference and further study. TEPCat is available at {\\tt http://www.astro.keele.ac.uk/$\\sim$jkt/tepcat/}} ",
        "introduction": "\\label{sec:intro}  Of the 550 planets known to orbit stars other than our Sun\\footnote{\\tt www.exoplanet.eu}, the transiting systems are the most interesting ones. Transiting extrasolar planets (TEPs) are the only planets outside our Solar system whose masses and radii, and thus surface gravities and mean densities, can be measured to reasonable precision. Observational selection effects mean that most known TEPs orbit very close to their parent star and thus have highly irradiated atmospheres whose physical properties can be scrutinised using high-quality astronomical observations.  One hindrance to the study of TEPs is that measurement of their physical properties requires not only transit light curves and radial velocity measurements of the parent stars, but also some sort of additional constraint. This is ordinarily obtained by forcing the properties of the star to match predictions from theoretical stellar models, guided by an effective temperature (\\Teff) and metal abundance (\\FeH) obtained from spectral analysis. The dependence on stellar theory leads to systematic errors which can be sizeable for some of the measured quantities, and also allows inhomogeneities to occur between studies which use different theoretical predictions or apply the constraint in a different way. This in turn compromises statistical studies of transiting planets.  For these reasons I am measuring the properties of the known TEPs using strictly homogeneous methods. Paper\\,I \\citep{Me08mn} discussed the methodology used to model the transit light curves, paying particular attention to error analysis and the treatment of limb darkening, and applied them to the 14 systems with good observational data at the time. Paper\\,II \\citep{Me09mn} explored the application of constraints from seven different sets of theoretical model predictions and alternatively an empirical mass--radius relation obtained from eclipsing binaries. Three of the sets of theoretical predictions were \\reff{selected to aid in the determination} of the physical properties of the same 14 systems, with detailed error budgets including random and systematic contributions. In Paper\\,III \\citep{Me10mn} I extended the number of systems to 30 and the number of theoretical predictions used to five sets, improving both the statistical weight of the ensemble and the precision of the systematic errorbars.  In the current work I enlarge the number of TEPs with homogeneous properties to \\reff{58}, concentrating on those which have been observed by the space missions \\kepler\\ \\citep{Borucki+10sci}, \\corot\\ \\citep{Baglin+06conf} and EPOCH \\citep{Christiansen+09iaus}. In Paper\\,II I outlined the concept of applying an observational mass--radius relation instead of using constraints from stellar theory, resulting in totally empirical measurements of the properties of TEP systems. This approach was not very successful (see Paper\\,III) because the mass-radius relation did not allow for stellar evolution and had to be calibrated on low-mass eclipsing binary systems whose properties (primarily radius) are affected by magnetic \\reff{activity} arising from comparatively fast rotation. In the current work I follow the alternative approach of \\citet{Enoch+10aa} which relates the radius of the star to its density, \\Teff\\ and \\FeH. This technique is not purely empirical, as it incorporates parameters derived using spectral synthesis techniques for both the TEP hosts and the calibrating sample, but returns results in much better agreement with the default method using theoretical predictions. Finally, I explore the opportunities for further study of the systems studied in this work.  For a small number of TEPs it is possible to either partially or totally avoid systematic errors: the study of transit timing variations (TTVs) allow the masses of some TEPs to be constrained directly \\citep{HolmanMurray05sci,Lissauer+11natur}; and radial velocities of the {\\em planet} HD\\,209458\\,b have been measured by \\citet{Snellen+10natur} from \\reff{infrared absorption} lines, allowing the system properties to be calculated in an identical way to double-lined eclipsing binary star systems. These methods are, however, not applicable to the predominant population of TEPs which show no detectable TTVs and are not amenable to direct velocity measurements. ",
        "conclusions": "\\label{sec:teps:summary}  The physical properties of 32 transiting extrasolar planetary systems have been derived from public light curves and published spectroscopic parameters of the host stars. These include 15 systems observed by the \\corot\\ satellite, ten by \\kepler, and five by the EPOCH project on the {\\it Deep Impact} spacecraft. Combined with the 30 objects examined in Paper\\,III, a sample of \\reff{58} TEPs with homogeneously measured properties is obtained.  All available transit light curves of each TEP were obtained and modelled using the {\\sc jktebop} code, with careful attention paid to the treatment of limb darkening, contaminating light, orbital eccentricity, Poisson and correlated noise, and long effective exposure times. The results for each light curve were then amalgamated to yield combined photometric parameters for the system, which were compared with literature results.  The physical properties of the TEPs were calculated from measured quantities by applying constraints from theoretical models, guided by the atmospheric parameters of the host stars. Five different sets of theoretical model tabulations were used, and the final results for each TEP are the unweighted mean of the individual results for each output parameter. Systematic errors were estimated by the interagreement between the individual model results, and statistical errors were propagated using a perturbation analysis. The constants and units needed in this process were tabulated for reference, and an error in the unit used for planetary density was fixed.  I also calculated the physical properties of each TEP system using a constraint obtained from eclipsing binary star systems (see also \\citealt{Enoch+10aa}). The constraint was applied in the form of $\\lten R = f(\\lten\\Teff, \\lten\\rho, \\MoH)$, where the precise equation and calibration coefficients were determined using the the measured properties of 90 well-studied detached eclipsing binaries. This gives results in generally good agreement with those from using theoretical stellar models as a constraint, although a trend towards poorer agreement is seen at higher metallicities. It is not obvious whether this trend arises from an imperfection in the calibration equation or source data, or from the physical effects included in the theoretical models.  The resulting physical properties of the 32 TEP systems are typically in good agreement with published results, but exceptions exist. My results for \\corot-5 disagree with those of the discovery paper, and this may be related to the treatment of orbital eccentricity. The public light curve of \\corot-8 does not match the published orbital ephemeris: I measure a revised ephemeris and somewhat different properties compared to the discovery paper. The two \\corot\\ light curves (short and long cadence) of \\corot-13 are discrepant. After rejection of the much less reliable 512\\,s cadence data I find physical properties of the system which are very different to previously thought. The resulting density of the planet is almost a factor of three smaller, moving it from outlier status to more representative of the general population of transiting Hot Jupiters. Many of the error estimates in the literature are far smaller than I find, and are not supported by the intrinsic quality of the data.  My analysis of the TEPs observed by \\kepler\\ uses data from Quarters 0, 1 and 2 for most of the objects. It is therefore the first analysis of most of the TEPs to include short-cadence data. This allows me to provide updated ephemerides and more reliable physical properties. My results for Kepler-5 are somewhat different to those previously published, due primarily to the inclusion of the Quarter 2 short-cadence data.  Asteroseismic studies are available for the three previously-known TEPs in the \\kepler\\ field, based on the \\kepler\\ short-cadence data, and for HD\\,17156 based on HST data. These studies use theoretical stellar models to interpret the oscillation spectrum of the star, and measure the stellar density to very high precision. The corresponding values I find from the light curve analysis are in good agreement for HD\\,17156 (0.8$\\sigma$) and TrES-2 (1.1$\\sigma$) but not for HAT-P-7 (2.9$\\sigma$) or HAT-P-11 (6.5$\\sigma$). This indicates a problem with at least one of the approaches, which might be related to underestimation of the true uncertainties, starspot activity or the measured orbital eccentricity of the HAT-P-11 system.  Finally, the complete error budgets generated for each TEP system allow identification of the observations which would lead to the greatest improvement in our measurement of their physical properties. Many objects would benefit from further photometric observations -- which continue to be obtained for the the TEPs in the \\kepler\\ field of view -- as well as from spectroscopic radial velocity measurements and spectral synthesis analyses. A list of nine TEP systems is given whose orbital ephemerides will become uncertain by more than one hour within this decade; the transits of \\corot-4 and \\corot-14 are already not predictable to within one hour.  The homogeneous physical properties obtained in this work will be useful for detailed statistical studies of the extrasolar planet population as well as for planning many types of follow-up observations of these objects. \\reff{The primary results from the current work and from previous papers in the series, along with a range of other useful information, have been concatenated and placed in an online catalogue. TEPCat is available at {\\tt http://www.astro.keele.ac.uk/$\\sim$jkt/tepcat/} in a range of convenient formats for readers to download for reference and further study.}"
    },
    "1107/1107.3837_arXiv.txt": {
        "abstract": "It has been realised only recently that TeV emission from blazars can significantly heat the intergalactic medium (IGM) by pair-producing high-energy electrons and positrons, which in turn excite vigorous plasma instabilities, leading to a local dissipation of the pairs' kinetic energy. In this work, we use cosmological hydrodynamical simulations to model the impact of this blazar heating on the Lyman-$\\alpha$ forest at intermediate redshifts ($z\\sim2-3$). We find that blazar heating produces an inverted temperature-density relation in the IGM and naturally resolves many of the problems present in previous simulations of the forest that included photoionisation heating alone. In particular, our simulations with blazar heating simultaneously reproduce the observed effective optical depth and temperature as a function of redshift, the observed probability distribution functions (PDFs) of the transmitted flux, and the observed flux power spectra, over the full redshift range $2<z<3$ analysed here. Additionally, by deblending the absorption features of Lyman-$\\alpha$ spectra into a sum of thermally broadened individual lines, we find superb agreement with the observed lower cutoff of the line-width distribution and abundances of neutral hydrogen column densities per unit redshift. Using the most recent constraints on the cosmic ultraviolet (UV) background, this excellent agreement with observations does not require rescaling the amplitude of the UV background; a procedure that was routinely used in the past to match the observed level of transmitted flux. We also show that our blazar-heated model matches the data better than standard simulations even when such a rescaling is allowed. This concordance between Lyman-$\\alpha$ data and simulation results, which are based on the most recent cosmological parameters, also suggests that the inclusion of blazar heating alleviates previous tensions on constraints for $\\sigma_8$ derived from Lyman-$\\alpha$ measurements and other cosmological data. Finally, we show that blazar heating dramatically alters the volume-weighted temperature PDF, implying an important change of the strengths of structure formation shocks (and thereby possibly particle acceleration in these shocks). The density PDF is also modified, suggesting that blazar heating may have interesting effects on structure formation, particularly on the smallest galaxies. ",
        "introduction": "\\label{sec:introduction}  The thermal history of the intergalactic medium (IGM) is of key importance for understanding cosmological structure formation since it determines the initial thermodynamic conditions for the formation of collapsed objects ranging from dwarf galaxies up to the scale of galaxy clusters. Simultaneously, the IGM acts as a cosmic heat reservoir that records reionisation processes in the Universe that inject substantial amounts of energy into the IGM on comparably short cosmological time scales. Detailed measurements of the thermal history of the IGM may therefore reveal the main contributors to the UV background, e.g., cosmic star formation or quasars, which are thought to be primarily responsible for these reionisation events.  Until recently, it was generally thought that all the main sources for heating the IGM had been identified and only the detailed interplay of these processes remained to be understood.  However, in a recent series of papers \\citep{Broderick2011,Chang2011,Pfrommer2011}, a novel heating process was proposed that taps efficiently into the energy reservoir of TeV blazars which are ultimately powered by accretion onto their central super-massive black holes (or more generally the engines of active galactic nuclei, AGN).  A subset of blazars emits the bulk of their energy in TeV $\\gamma$-ray radiation ($E\\gtrsim 100~\\rmn{GeV}$) to which the Universe is opaque; i.e., these energetic gamma rays necessarily annihilate upon the photons comprising the extragalactic background light (EBL), producing an ultrarelativistic population of electron-positron pairs. Typical gamma-ray mean free paths range from 30 Mpc to 1 Gpc, depending upon $\\gamma$-ray energy and source redshift. The pairs produced by TeV $\\gamma$-rays have typical Lorentz factors of $10^5-10^7$. Despite its dilute nature, a beam of these ultra-relativistic pairs is susceptible to plasma beam instabilities while propagating through the IGM. In particular, the ``oblique'' instability, a more virulent and robust cousin of the commonly discussed Weibel and two-stream instabilities, appears to grow out of the pair beam anisotropy as a source of free energy \\citep{Bret-Firp-Deut:04,Bret-Firp-Deut:05,Bret:09,Bret-Grem-Diec:10,Lemo-Pell:10}. Though the nonlinear evolution of these instabilities is unknown, they may saturate at a level such that they dissipate the kinetic energy of the beam at a rate comparable to their linear growth rate, thereby heating the IGM locally, following the suggestion of \\citet{Broderick2011}.  If this scenario or an analogous, similarly efficient mechanism operates in practise, it necessarily suppresses the inverse Compton cooling of the pairs, which acts on a longer timescale. This naturally explains the absence of an inverse Compton bump in observed TeV-blazar spectra without invoking an intergalactic magnetic field \\citep[][and references therein]{Broderick2011}. At the same time, it allows for a redshift evolution of TeV blazars that is identical to that of quasars without overproducing the extragalactic $\\gamma$-ray background (EGRB).  In fact, for plausible parameters of TeV-blazar spectra, it is possible to explain the high-energy part of the EGRB \\citep{Broderick2011}.  This is in stark contrast to the previous picture which assumed that the TeV $\\gamma$-ray emission is reprocessed into the GeV regime by means of inverse Compton processes. In this picture, in order to not overproduce the EGRB observed by {\\em Fermi} \\citep{Fermi_EGRBApJ2010}, the spectral redistribution of TeV $\\gamma$-ray energy into the GeV regime strongly constrains the evolution of the luminosity density of TeV blazars. In these studies it was found that TeV blazars cannot exhibit the dramatic rise in number seen in the quasar distribution by $z\\sim1$--$2$, i.e., the comoving number of blazars must have remained essentially fixed \\citep[see, e.g.,][]{Naru-Tota:06,Knei-Mann:08,Inou-Tota:09,Vent:10}; at odds with both the large-scale mass assembly history of the Universe as traced by the star formation history, and with the luminosity history of similarly accreting systems, e.g., the quasar luminosity density.  Integrating over the energy flux per mean free path of all known TeV blazars yields a luminosity density, or equivalently a local heating rate, that dominates that of photoheating by more than an order of magnitude at the present epoch at mean density, after accounting for incompleteness corrections \\citep{Chang2011}. As demonstrated in \\citet{Broderick2011}, the local TeV blazar luminosity function is consistent with a scaled version of the quasar luminosity function \\citep{Hopkins+07}, thus the conservative assumption is that they evolve similarly, presumably due to the same underlying accretion physics.  With this assumption, blazar heating starts to become important around the epoch of He {\\sc ii} reionisation ($z\\sim3.5$). By $z=0$, the blazar heating rate is more than an order of magnitude larger in comparison with the photoheating rate alone.  Because the Universe is opaque to the propagation of TeV photons over cosmological distances, the redshift evolution of the TeV blazar luminosity density is observationally inaccessible for $z\\gtrsim 0.5$. Also, the exact value of the blazar heating rate is subject to an uncertain incompleteness correction factor, which accounts for the incomplete sky coverage of current TeV instruments, the duty cycle of TeV blazars, their spectral variability, variations of the TeV-blazar redshift evolution, and additional TeV sources that may contribute to the plasma instability heating. Hence we can try to invert the problem and investigate how we can use the IGM as a cosmic calorimeter to probe the physics of TeV $\\gamma$-rays and the global properties of TeV blazars by means of the energy that is deposited through plasma instabilities (or similarly efficient dissipation processes of the TeV emission from blazars).  In this manner, the thermal history of the IGM provides an analogous argument to that of \\citet{Solt:82}, yielding an important constraint on the redshift evolution of the blazar luminosity density \\citep{Chang2011}. Hence we turn to precision measurements of the Lyman-$\\alpha$ forest at $z\\sim 2-3$ \\citep[e.g.,][]{Viel2009} in this work, with the goal to clarify the possible signatures of such an additional heating process, which is characterised by some unusual properties as we will discuss now.  Since the number densities of EBL photons and of TeV blazars are nearly homogeneous on cosmological scales, so is the resulting pair density. Hence, the implied heating rate is also homogeneous, i.e., the volumetric blazar heating rate is expected to be spatially uniform and {\\em independent} of IGM density. This implies that the energy deposited per baryon is substantially larger in more tenuous parts of the Universe, such that underdense regions experience a larger temperature increase, producing an inverted temperature-density (T-$\\rho$) relation with an asymptotic scaling of $T\\propto 1/\\rho$ for densities much lower than the cosmic mean. Interestingly, detailed studies of high-resolution Lyman-$\\alpha$ forest spectra led to the conjecture that such an inverted T-$\\rho$ relation in fact exists at $z=2-3$ \\citep{Bolton2008,Viel2009}, something which has proven hard to arrange within the context of standard reionisation models \\citep{McQuinn2009,Bolton2009}. Related to that, when deblending the absorption features of Lyman-$\\alpha$ spectra into a sum of thermally broadened individual lines and measuring their strengths and widths, simulations produce line widths that are significantly narrower than observed, implying an associated inconsistency between the simulated and measured line width distribution \\citep{Bryan1999,Bryan2000,Theuns1999,Machacek2000}. This finding is largely independent of the amount of small-scale power, but does depend strongly on the IGM temperature; heuristically increasing the temperature of the IGM can hence reproduce the observed line-width distribution rather well \\citep{Theuns2000}.  In this work, we perform state-of-the-art simulations of the Lyman-$\\alpha$ forest that include different variants of blazar heating (to address the uncertainty of the overall heating rate), and we compare the predictions to data in observational space. In this way, we obtain important information on the validity of the blazar heating scenario and we can elucidate whether it helps to resolve current problems in understanding some aspects of the IGM observations.  After describing our methods for simulating the Lyman-$\\alpha$ forest in Sect.~\\ref{sec:methods}, we present our results in Sect.~\\ref{sec:results}. We focus on the thermal history and the T-$\\rho$ relation of the IGM in Sect.~\\ref{sec:th_eos_igm}, the effective optical depths and photoionisation rates in Sect.~\\ref{sec:optical_depths}, the transmitted flux probability distribution functions and power spectra in Sect.~\\ref{sec:flux_pdfs_power_spec}, and the Voigt profile fitting of the Lyman-$\\alpha$ forest in Sect.~\\ref{sec:line_fitting}. Finally, we give an outlook on implications for cosmological parameter estimation when including the physics presented in this work in Sect.~\\ref{sec:outlook}, and conclude in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Recent studies discovered that TeV $\\gamma$-ray emission from blazars can significantly heat the low-density IGM by annihilating and pair producing on the extragalactic background light.  The resulting high-energy electrons and positrons excite plasma beam instabilities which dissipate the pairs' energy locally \\citep{Broderick2011,Chang2011}, leading to important implications for structure formation \\citep{Pfrommer2011}. In this picture, the IGM acts as a cosmic heat bath that absorbs most of the bolometric energy of TeV blazars, which are ultimately powered by accretion onto their central super-massive black holes. Since the distributions of EBL photons and of TeV blazars are nearly homogeneous on cosmological scales, the resulting pair density is spatially homogeneous as well. As a result, the implied heating rate is independent of the IGM density.  We have included the expected heating term in state-of-the-art cosmological hydrodynamical simulations and investigated its impact on the thermal history of the IGM and the properties of the Lyman-$\\alpha$ forest. We find that:  \\begin{itemize}  \\item In agreement with linear theory predictions by \\citet{Chang2011}, the temperature of the low density IGM is strongly boosted by blazar heating, resulting in an inverted temperature-density relation as favoured by Lyman-$\\alpha$ forest observations \\citep{Viel2009}.  \\item At $z<3$, when blazar heating is the dominant heating source, the IGM temperature, as a function of redshift, is also in good agreement with current constraints based on the curvature of observed Lyman-$\\alpha$ forest spectra \\citep{Becker2011}.  \\item The mean transmissions predicted by runs with blazar heating are in excellent agreement with observations when employing the photoionisation rates predicted by up-to-date models of the cosmic UV background (FG'09). We adopt these rates in our {\\em self-consistent} analysis which uses the same UV background for simulations and analysis. To allow for the uncertainties in the determination of the photoionisation rates, we additionally explore how results change if we {\\em tune} the UV background to match the observed mean transmission values.  \\item Blazar heating alters the distribution functions and power spectra of the transmitted flux and resolves discrepancies with observational constraints that exist when photoionisation is the only modelled source of heating. This improved match of simulations and data in our blazar heating models is also true if we rescale the photoionisation background to match the observed mean transmission at each redshift. This seems to be a unique property of the inverted $T$--$\\rho$ relation \\citep{Bolton2008} which is naturally produced by blazar heating.  \\item Heating of the IGM by blazars shifts the lower cutoff of Lyman-$\\alpha$ absorption line widths to larger values due to increased thermal broadening and decreases the abundance of weak lines with column densities $N_\\rmn{HI}\\lesssim3\\times 10^{13} \\,\\rmn{cm}^{-2}$. This makes the lower cutoff and the abundance of lines per unit redshift and column densities consistent with observations and provides additional evidence for high IGM temperatures independent of whether or not we rescale to the observed mean transmission.  \\item The inclusion of blazar heating brings the simulated large-scale flux power into agreement with the observed level. The simulation predictions are based on the most recent cosmological parameters which suggests that previous tensions in Lyman-$\\alpha$ measurements of $\\sigma_8$ when compared to other cosmological constraints are alleviated by this novel heating mechanism. Our results also support previous claims that the empirical assumption of an inverted temperature-density relation for the IGM already relieves the tension in the constraints on $\\sigma_8$ \\citep{Viel2009}.  \\item Blazar heating dramatically alters the volume-weighted temperature PDF which implies that the strengths of first structure formation shocks during non-linear collapse are weakened. Future work is needed to work out the implications for the particle acceleration efficiencies and the resulting non-thermal emission from galaxy clusters. Modifications in the density PDF furthermore suggest that blazar heating may have interesting effects on structure formation, particularly on the scale of dwarf galaxies.  \\end{itemize}  To summarise, including blazar heating in cosmological simulations results in properties of the IGM and the Lyman-$\\alpha$ forest that are in excellent agreement with observations. In particular, we demonstrate in this work that our intermediate blazar heating model in combination with the FG'09 UV background represents the first successful, physically-motivated model that simultaneously reproduces all Lyman-$\\alpha$ forest data (flux PDFs, power spectra, temperature evolution, line width and column density distributions) without the need to rescale the UV background at each redshift. This strengthens the case that this novel heating mechanism is indeed operating in our Universe, and at the same time it resolves long standing difficulties in our understanding of the Lyman-$\\alpha$ forest.  This impressive improvement in modelling the Lyman-$\\alpha$ forest appears to be a direct consequence of the peculiar properties of blazar heating. Its heating rate, which is independent of IGM density, naturally produces the inverted temperature-density relation that Lyman-$\\alpha$ forest data suggest \\citep{Viel2009}. The recent and continuous nature of the heating generates the redshift evolutions of the IGM temperature, the effective optical depth, the flux PDF, and flux power spectrum that accurately match observed trends for the redshift interval $2<z<3$ analysed here. Finally, the magnitude of the heating rate required by Lyman-$\\alpha$ forest data approximately coincides with the total energy output of TeV blazars (or equivalently $\\sim0.2\\%$ of that of quasars).  While all these properties are quite unusual in the context of H and He {\\sc i/ii} photoheating and traditional feedback processes, e.g., that of AGN or supernova driven winds, they are natural in the TeV-blazar heating mechanism. There is hence ample motivation for further studies of this scenario, which are also needed to establish in full detail how blazar heating affects cosmological constraints and structure formation over a range of scales."
    },
    "1107/1107.1373_arXiv.txt": {
        "abstract": "For the first time, we investigate the X-ray/infrared (IR) correlation for star-forming galaxies at $z \\sim$\\,1, using SPIRE submm data from the recently-launched \\textit{Herschel Space Observatory} and deep X-ray data from the 2\\,Ms \\textit{Chandra} deep field north (CDFN) survey. We examine the X-ray/IR correlation in the soft X-ray (SX, 0.5--2\\,keV) and hard X-ray (HX, 2--10\\,keV) bands by comparing our $z \\sim$\\,1 SPIRE-detected star-forming galaxies (SFGs) to equivalently IR-luminous ($L_{\\rm IR}$\\,$>$\\,10$^{10}$\\,L$_{\\odot}$) samples in the local/low redshift Universe. Our results suggest that the X-ray/IR properties of the SPIRE SFGs are on average similar to those of their local counterparts, as we find no evidence for evolution in the $L_{\\rm SX}$/$L_{\\rm IR}$ and  $L_{\\rm HX}$/$L_{\\rm IR}$ ratios with redshift. We note however, that at all redshifts, both $L_{\\rm SX}$/$L_{\\rm IR}$ and  $L_{\\rm HX}$/$L_{\\rm IR}$ are strongly dependent on IR luminosity, with luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs, $L_{\\rm IR}$\\,$>$\\,10$^{11}$\\,L$_{\\odot}$) having up to an order of magnitude lower values than normal infrared galaxies ($L_{\\rm IR}$\\,$<$\\,10$^{11}$\\,L$_{\\odot}$). We derive a $L_{\\rm SX}$-$L_{\\rm IR}$ relation and confirm the applicability of an existing $L_{\\rm HX}$-$L_{\\rm IR}$ relation for both local and distant LIRGs and ULIRGs, consistent with a scenario where X-ray luminosity is correlated with the star-formation rate (SFR). ",
        "introduction": "For at least two decades, the infrared/X-ray connection has been an important point of focus in extragalactic surveys. The various X-ray missions (\\textit{Einstein}, \\textit{ROSAT}, \\textit{ASCA}, \\textit{BeppoSAX}, \\textit{Chandra} and \\textit{XMM-Newton}) together with the infrared observatories (\\textit{IRAS}, \\textit{ISO} and \\textit{Spitzer}) have established the X-ray/infrared synergy not only as a tool for studying the active galactic nucleus/host galaxy interplay, but also star-formation and stellar evolution (a few examples of the multitude of papers on this topic: Eales $\\&$ Arnaud 1988\\nocite{EA88}; Fabbiano 1988\\nocite{Fabbiano88}; Fabbiano, Gioia $\\&$ Trinchieri 1988\\nocite{FGT88}; Rieke 1988\\nocite{Rieke88}; Green et al. 1989\\nocite{Green89}; Rephaeli et al. 1991\\nocite{Rephaeli91}; Green et al. 1992\\nocite{Green92}; Boller et al. 1992\\nocite{Boller92}; Barcons et al. 1995\\nocite{Barcons95}; Rephaeli, Gruber $\\&$ Persic 1995\\nocite{RGP95}; Imanishi $\\&$ Ueno 1999\\nocite{IU99}; Severgnini et al. 2000\\nocite{Severgnini00}; Risaliti et al. 2000\\nocite{Risaliti00}; Cavaliere, Giacconi $\\&$ Menci 2000\\nocite{CGM00}; Fadda et al. 2002\\nocite{Fadda02}; Gandhi $\\&$ Fabian 2003\\nocite{GF03}; Manners et al. 2004\\nocite{Manners04}; Panessa et al. 2005\\nocite{Panessa05}; Netzer et al. 2005\\nocite{Netzer05}; Franceschini et al. 2005\\nocite{Franceschini05}; Hickox et al. 2007\\nocite{Hickox07}; Treister et al. 2009\\nocite{Treister09}; Trichas et al. 2009\\nocite{Trichas09}; Lutz et al. 2010\\nocite{Lutz10}; Park et al. 2010\\nocite{Park10}; Goulding et al. 2011\\nocite{Goulding11}).   Emission in the infrared (IR) to millimetre ($\\sim$\\,8--1000\\,$\\mu$m) part of the spectrum is attributed to galactic dust being heated to a range of temperatures by incident UV and optical photons. In cases where an active galactic nucleus (AGN) is present, emission from the torus (or other axisymmetric dust/gas configurations, e.g. Efstathiou $\\&$ Rowan-Robinson 1995\\nocite{ERR95}) can also play a prominent role at short infrared wavelengths ($<$\\,20\\,$\\mu$m, rest-frame), as torus dust is heated to near-sublimation temperatures by the AGN radiation field. On the other hand, flux longwards of $\\sim$\\,20\\,$\\mu$m is strongly correlated with the star formation rate (SFR), although in the case of quiescently star-forming galaxies, emission from dust heated by older stellar populations might also contribute significantly to the far-IR (e.g. Bendo et al. 2010\\nocite{Bendo10}, Calzetti et al. 2010\\nocite{Calzetti10}). Nevertheless, in galaxies with high total infrared luminosities ($L_{\\rm IR}$\\,$>$10$^{10}$\\,L$_{\\odot}$) where large dust masses are implied, (i) the source of heating is intense and (ii) the dust replenishment is continuous, making IR emission an excellent tracer of young star-forming/starburst regions. Moreover, contamination from an AGN (if present) is minimal in the far-IR/submm part of the spectrum. Recently, Hatziminaoglou et al. (2010\\nocite{Hatziminaoglou10}) showed that the submm colours of AGN hosts are indistinguishable from those of the general submm-detected population.  Emission in the high energy part of the spectrum is a combination of both AGN (if present) and star-formation related processes. AGN are the most powerful X-ray emitters, with X-rays originating in the hot corona surrounding the accretion disk through inverse Compton scattering of softer energy photons (Sunyaev $\\&$ Titarchuk 1980\\nocite{ST80}), generally amounting to a few per cent of the AGN energy budget (Elvis 1994\\nocite{Elvis94}). On the other hand, the applicability of X-rays as a star formation tracer is linked to a galaxy's star formation history and interstellar medium (ISM) conditions (e.g. Ghosh $\\&$ White 2001\\nocite{GW01}; Fabbiano et al. 2004\\nocite{Fabbiano04}). Star-forming galaxies have complex 0.5--10\\,keV spectra, with contributions from a low temperature soft component (k$T\\lesssim$1\\,keV) and a harder component which is either predominantly thermal (k$T\\sim$5--10\\,keV) or non-thermal (power-law index $\\Gamma$\\,$\\sim$1.5--2) or a combination of both (Persic $\\&$ Rephaeli 2002\\nocite{PR02}). The soft component is primarily thermal emission from gas in the ISM heated to X-ray temperatures by stellar winds and supernovae, with a contribution from point sources of the order of 30 per cent, according to Chandra observations of local galaxies (e.g. Fabbiano, Zezas $\\&$ Murray 2001\\nocite{FZM01}; Zezas, Ward $\\&$ Murray 2003\\nocite{ZWM03}). On the other hand, more than 60 per cent of the hard component is resolved into point sources and associated with X-ray binaries (e.g. Griffiths et al. 2000\\nocite{Griffiths00}).  High mass X-ray binary (HMXB) spectra can be represented by a power law up to about 20\\,keV, whereas the spectra of low mass X-ray binaries (LMXB) are a combination of thermal bremmstrahlung and black-body emission in the $<$20\\,keV region (e.g. Christian $\\&$ Swank 1997\\nocite{CS97}, see also Persic $\\&$ Rephaeli 2002). HMXBs, where the companion star is massive and short-lived, are direct tracers of recent star-formation, whereas LMXBs have lifetimes of the order of 10\\,Gyr and are hence more appropriate tracers of stellar mass (e.g. Ptak et al. 2001\\nocite{Ptak01}; Grimm, Gilfanov $\\&$ Sunyaev 2002, 2003\\nocite{GGS02}\\nocite{GGS03}). Although the relative timescales of X-ray emitting processes play an important role in how one interprets the star-formation history of quiescently star-forming galaxies or starburst (SB) systems, X-rays are a good overall tracer of the first $\\sim$\\,30\\,Myr of star-forming activity (e.g. Mass-Hesse, Oti-Floranes $\\&$ Cervino 2008\\nocite{MOC08}), with the relation between star formation rate (SFR) and full band X-ray luminosity found to be linear in the $>$1\\,M$_{\\odot}\\,\\rm yr^{-1}$ SFR regime (Gilfanov, Grimm $\\&$ Sunyaev 2004\\nocite{GGS04}).   As both X-ray and IR emission trace star formation, one would expect these quantities to be correlated in the absence of an AGN. The X-ray/IR correlation, described as log\\,($L_{\\rm  X}$)=A\\,+\\,B\\,log\\,($L_{\\rm IR}$), has been studied extensively in both the soft (0.5--2\\,keV) and hard (2--10\\,keV) energy bands with samples of local/low-redshift galaxies (eg. Griffiths $\\&$ Padovani 1990\\nocite{GP90}; David, Jones $\\&$ Forman 1992\\nocite{DJF92}; Ranalli, Comastri $\\&$ Setti 2003\\nocite{RCS03}; Franceschini et al. 2003\\nocite{Franceschini03b}; Rosa-Gonz{\\'a}lez et al. 2007). These studies have calculated different values for the slope of the soft X-ray/IR relation with B roughly ranging from 0.7 to 1. A linear relation (B=1) is a strong indication that X-ray and IR emission directly trace the star formation rate, however deviations from linearity do not necessarily suggest the contrary. The large range in computed slopes could be a consequence of the small number statistics in the samples studied, in combination with the large scatter in the galaxies' physical properties, such as the efficiency of energy coupling to the ISM which heats the X-ray emitting gas, varying contributions from point sources, as well as the amount of absorption undergone by the soft X-rays due to gas/dust columns in the line of sight. In that respect, hard X-rays are a more reliable tracer of star-formation, as the ISM is mostly transparent at energies above \\,$\\sim$\\,2keV, except for dense, massive molecular cloud environments (Grimm, Gilfanov $\\&$ Sunyaev 2002; Ranalli, Comastri $\\&$ Setti 2003; Persic et al. 2004\\nocite{Persic04}). Indeed, hard X-ray emission linearly traces the star formation rate (e.g. Grimm, Gilfanov $\\&$ Sunyaev 2003; Persic et al. 2004), although because of the lower sensitivity of detectors in the hard X-rays, this band is still not as extensively used. In addition, whereas soft X-rays suffer little contamination from X-ray binary emission, to use hard X-rays successfully, the relative contribution of HMXBs and LMXBs needs to be quantified. This is easier for SB environments where HMXBs dominate but for more quiescently star-forming systems, LMXBs, which trace the stellar mass, can contribute substantially, introducing non-linearities (e.g. Grimm, Gilfanov $\\&$ Sunyaev 2002; Persic et al 2004\\nocite{Persic04}).  So, what is the role of the X-ray/IR correlation in extragalactic surveys? Apart from the obvious motivation, i.e. examining the physical properties of different galaxy populations, one of its key applications could be to discriminate between AGN-dominated and star-formation dominated systems. The ability to discriminate between the two is particularly significant in X-ray surveys, where AGN dominate the bright end of the number counts, but the star-forming galaxy population is expected to outnumber AGN at faint fluxes (f$_{\\rm 0.5\u2013-2\\,keV}$\\,$\\sim$\\,10$^{-17}$\\,erg\\,s$^{-1}$\\,cm$^{-2}$; e.g. Hornschemeier et al. 2002\\nocite{Hornschemeier02}, 2003\\nocite{Hornschemeier03}; Bauer et al. 2004; Brandt $\\&$ Hasinger 2005\\nocite{BH05}; see also Georgakakis et al. 2003\\nocite{Georgakakis03}; 2004\\nocite{Georgakakis04}; 2006ab\\nocite{Georgakakis06a}\\nocite{Georgakakis06b}; 2007\\nocite{Georgakakis07}; 2008\\nocite{Georgakakis08}). As previously shown (e.g. Georgakakis et al. 2007\\nocite{Georgakakis07}), the X-ray/IR correlation is a useful way of separating starbursts from AGN. However, up to now, such studies have used relations applicable to normal IR galaxies in the local Universe. What remains to be settled is whether there is a change in the characteristics of the X-ray/IR relation with redshift and/or SFR. As noted earlier, a linear relationship for star-forming galaxies is a strong indication that X-ray and IR emission directly trace the star formation rate. As a result, determining the X-ray/IR relation would play an important role in calibrating X-ray emission as a star formation indicator, providing an alternative census of the global SFR density.  Although various extragalactic X-ray studies have targetted cosmologically significant redshifts (e.g. Norman et al. 2004\\nocite{Norman04}; Ptak et al. 2007\\nocite{Ptak07}; Georgakakis et al 2007; Lehmer et al. 2008\\nocite{Lehmer08}), the X-ray/IR correlation has been examined only for local galaxies (e.g. Ranalli, Comastri $\\&$ Setti et al. 2003) and X-ray selected samples at moderate redshifts (Grimm, Gilvanov $\\&$ Sunyaev 2002; Rosa-Gonz{\\'a}lez et al. 2007\\nocite{Rosa-Gonzalez07}), mainly due to sensitivity limitations in X-ray surveys. For the first time, we are able to investigate the X-ray/IR correlation at $z$\\,$\\sim$1, for submm-bright sources with high star-formation rates, using deep X-ray data from the \\textit{Chandra} 2\\,Ms CDFN survey and sub-mm data from SPIRE (Griffin et al. 2010\\nocite{Griffin10}), the sub-mm bolometer array on the recently-launched \\textit{Herschel Space Observatory}\\footnote{\\textit{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. 2010\\nocite{Pilbratt10}). Our aim is to compare the X-ray properties of $z$\\,$\\sim$1 IR-luminous galaxies to their counterparts in the local ($z<0.1$), moderate ($z\\sim0.6$) and high ($z\\sim2$) redshift Universe. The paper is laid out as follows: in Section \\ref{sec:sample} we present the sample and in Sections \\ref{sec:soft} and \\ref{sec:hard} we examine the soft and hard X-ray/IR relations. Our summary and conclusions are presented in Section \\ref{sec:conclusions}. Throughout this paper, we adopt a concordance consmology of H$_0$=70\\,km\\,s$^{-1}$Mpc$^{-1}$, $\\Omega_{\\rm M}$=1-$\\Omega_{\\rm \\Lambda}$=0.3. ",
        "conclusions": "\\label{sec:conclusions}  Our analysis is based on a sample of 67, spectroscopically-identified, SPIRE-detected, submm-bright ($f_{250 \\rm \\mu m}$\\,$>$17.4\\,mJy; $>$3$\\sigma$ including confusion noise) galaxies in the GOODS-N field, observed as part of the \\textit{Herschel}/HerMES survey, the majority of which are LIRGs and ULIRGs at z\\,$\\sim$1. Taking advantage of the 2\\,Ms \\textit{Chandra} Deep Field North survey, we investigate the SPIRE sources' soft and hard X-ray emission relative to their infrared output and compare our results to several other samples from the literature. The focus of our study is the X-ray/IR correlation for star-forming IR-luminous galaxies at z$\\sim$1. Our sample is selected to be submm-bright, implying that all objects are star-formation-dominated in the infrared, thus we identify AGN through X-ray diagnostics and optical spectral lines. We class as AGN-dominated in the X-rays, objects which satisfy at least one of the following criteria: a hard X-ray/R-mag flux ratio (log\\,[f$_{\\rm HX}$/f$_{\\rm R}$]) greater than -1, a hardness ratio (HR) greater than -0.1 and the presence of high-excitation or broad optical lines. Out of the 43 X-ray detected sources, 18 are identified as AGN, amounting to a fraction of 27 ($\\pm$10) per cent in the SPIRE sample. Note that our aim is to exclude any obvious AGN, but we cannot ensure a pure star-forming sample, since low luminosity and/or unabsorbed AGN will be missed. Once obvious AGN are eliminated, our analysis is based on 24 X-ray detected SPIRE objects and 23 non-X-ray detected SPIRE sources, whose X-ray properties we assess through stacking.  We conclude that the X-ray/IR properties of star-formation-dominated SPIRE galaxies, quantified by the soft and hard X-ray/IR luminosity ratios ($L_{\\rm SX}$/$L_{\\rm IR}$ and $L_{\\rm HX}$/$L_{\\rm IR}$), are consistent with those observed in the local ($z<0.1$) Universe. They are also consistent with the ratios of equivalent sources at more moderate ($z\\sim0.6$) and higher redshifts ($z\\sim2$), indicating that, overall, there is no evidence for evolution of the X-ray/IR correlation for IR-luminous galaxies with redshift. Note that, although none of the samples we present here are entirely impervious to AGN contamination, the main outcome of our study is robust, in the sense that once all identifiable bona-fide AGN are excluded, the remaining sources have consistent average X-ray/IR ratios at all redshifts.  The linear nature of the X-ray/IR correlation describes a scenario whereby both X-ray and IR emission trace the star-formation rate and is thus important in calibrating X-ray emission as a star formation indicator, providing an alternative census of the global SFR density. This will be particularly applicable in future deeper X-ray surveys, where the X-ray number counts at the faint flux end will be dominated by star-forming galaxies.     \\onecolumn \\begin{table} \\centering \\caption{The full SPIRE sample of 84 SPIRE-detected, submm-bright ($f_{250 \\rm \\mu m}$\\,$>$17.4\\,mJy) galaxies in the GOODS-N field, observed as part of the \\textit{Herschel}/HerMES survey. The columns are as follows: (1) ID, (2) (3) IRAC-3.6\\,$\\mu$m RA and DEC, (4) spectroscopic redshift (available only for 67 sources), (5) total infrared luminosity (8--1000\\,$\\mu$m), (6) the cross-matched ID from the IC catalogue of HDFN Chandra data, (7) (8) (9) (10) the flux and luminosity in the soft and hard bands (SX stands for `soft X-ray band' and HX stands for `hard X-ray band'), (11) the hardness ratio (HR), (12) the $f_{\\rm HX}$/$f_{\\rm R}$ ratio used to identify AGN (see section \\ref{sec:agn_fraction}), computed for sources which are detected in both the R-band and the hard X-rays. Note that, although we provide the corresponding HDFN X-ray ID from the IC catalogue, the fluxes and upper limits we present here are all derived using the IRAC-3.6\\,$\\mu$m position of each SPIRE source to extract counts from the X-ray map down to a Poisson significance $<3\\times10^{-3}$ --- see Section \\ref{sec:xraydata} for details. `$<$\\,f$_{\\rm X}$' corresponds to the 3$\\sigma$ upper limit. Column (13) refers to the classification of the source by optical spectroscopy (Barger et al. 2005; Trouille et al. 2008). SFG stands for `Star Forming Galaxy', sources with strong Balmer lines and no broad or high-ionization lines. HEL AGN stands for `High Excitation Line AGN', sources with [NeV] or CIV lines or strong [OIII] lines, i.e. EW([OIII] $\\lambda$\\,5007) $>$ 3EW(H$\\beta$). BL AGN stands for `Broad-line AGN', sources with FWHM line widths greater than 2000\\,km\\,s$^{-1}$. Note that the Trouille et al. (2008) classification applies only to X-ray sources. Where there is no information available, the spectrum could not be classified. Column (14) refers to whether the source is SB or AGN dominated in the X-rays, according to the criteria outlined in section \\ref{sec:agn_fraction}. The 3 ambiguous sources (labelled `Amb'), are included in the SFG sub-sample, but our conclusions do not change if instead they are included in the AGN sub-sample. } \\begin{tabular}{|@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l|}  \\hline \\hline ID&RA& Dec. & z & $L_{\\rm IR}$ & HDFN ID & $f_{\\rm SX}$   & $f_{\\rm HX}$   & $L_{\\rm SX}$ & $L_{\\rm HX}$ & HR & $f_{\\rm HX}$/$f_{\\rm R}$ &Spectral & X-ray\\\\ & deg & deg &   & log\\,(L$_\\odot$) &         & erg/s/cm$^{2}$ & erg/s/cm$^{2}$ & log (erg/s)     & log (erg/s)     &    &  & Class    & Class\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)&(14)\\\\ \\hline 1&189.28772&62.38866&1.26&12.39&---                &$<$2.37e-16&$<$7.67e-16&---&---&---&--- &---     &---\\\\ 2&189.27887&62.38520&0.50&11.42&---                &$<$1.39e-16&$<$2.98e-16&---&---&---&--- &---     &---\\\\ 3&189.29427&62.37627&1.52&12.39&---                &$<$1.34e-16&$<$3.39e-16&---&---&---&--- &---     &---\\\\ 4&189.27174&62.33079&0.84&11.75&---                &$<$7.65e-17&$<$8.15e-16&---&---&---&--- &---     &---\\\\ 5&189.35646&62.32817&0.28&10.75&hdfn-358            &1.86e-16&$<$3.93e-16&40.64&---&-1&---  &SFG      &SFG\\\\ 6&189.14424&62.32388&0.87&11.72&hdfn-086            &1.62e-16&9.04e-15&41.75&43.50&0.83$^{0.86}_{0.81}$&0.60$^{+0.03}_{-0.03}$&SFG      &AGN\\\\ 7&189.08646&62.31671&---&---&---                &$<$5.26e-17&$<$2.79e-16&---&---&---&---&---      &---\\\\ 8&189.54642&62.30510&---&---&---                &$<$2.44e-16&$<$7.17e-16&---&---&---&---   &---   &---\\\\ 9&189.22009&62.30213&0.25&10.90&---                &2.05e-17&$<$1.94e-16&39.59&---&-1&---   &---   &SFG\\\\ 10&189.44516&62.29540&1.15&11.81&---                &$<$1.11e-16&$<$3.15e-16&---&---&---&---   &---   &---\\\\ 11&189.39382&62.28980&0.64&11.62&---                &4.24e-17&$<$1.24e-16&40.85&---&-1&---  &---    &SFG\\\\ 12&189.51512&62.28651&0.28&11.01&---                &$<$1.55e-16&$<$3.56e-16&---&---&---&---   &---   &---\\\\ 13&189.08741&62.28600&1.99&12.63&---                &$<$7.07e-17&$<$3.02e-16&---&---&---&---  &---    &---\\\\ 14&189.57885&62.28439&---&---&---                &$<$7.20e-17&$<$6.14e-16&---&---&---&---  &---    &---\\\\ 15&189.13544&62.28315&0.44&11.26&---               &$<$2.61e-17&$<$2.75e-16&---&---&---&---   &---   &---\\\\ 16&189.31910&62.27872&0.56&11.28&hdfn-264            &4.54e-17&1.95e-16&40.73&41.37&0.03$^{+0.28}_{-0.20}$&-1.18$^{+0.22}_{-0.78}$&HEL AGN  &AGN\\\\ 17&189.37984&62.27229&0.98&11.89&---               &$<$7.81e-17&$<$1.31e-16&---&---&---&---   &---   &---\\\\ 18&188.99886&62.26384&0.38&11.38&hdfn-335            &1.55e-16&$<$4.42e-16&40.86&---&-1&--- &SFG      &SFG\\\\ 19&189.07638&62.26402&2.00&12.70&hdfn-291            &3.01e-17&$<$3.87e-16&41.89&---&-1&--- &---     &SFG\\\\ 20&188.99144&62.26019&0.38&11.60&hdfn-503            &4.66e-17&$<$4.47e-16&40.34&---&-1&--- &---     &SFG\\\\ 21&189.06708&62.25379&2.58&12.91&hdfn-117            &2.01e-16&1.12e-15&42.98&43.72&0.24$^{+0.08}_{-0.07}$&0.69$^{+0.09}_{-0.10}$&---     &AGN\\\\ 22&189.48220&62.25186&0.19&10.53&hdfn-356            &4.61e-16&5.48e-18&40.66&38.74&-0.59$^{+0.13}_{-0.14}$&-3.85$^{+1.26}_{-0.56}$&---     &SFG\\\\ 23&188.96211&62.24923&0.39&11.03&---                &$<$2.65e-16&$<$1.41e-15&---&---&---&---   &---   &---\\\\ 24&189.14271&62.24249&0.52&11.20&---                &$<$2.49e-17&$<$1.77e-16&---&---&---&---    &---  &---\\\\ 25&189.14832&62.23999&2.01&12.88&hdfn-082            &1.56e-16&1.73e-15&42.61&43.65&0.37$^{+0.07}_{-0.07}$&0.25$^{+0.05}_{-0.05}$&HEL AGN  &AGN\\\\ 26&189.53851&62.23786&---&---&---                &$<$1.08e-16&$<$7.02e-16&---&---&---&---     &---&--- \\\\ 27&189.26141&62.23376&1.25&12.03&---               &$<$3.21e-17&$<$1.14e-16&---&---&---&---     &---&--- \\\\ 28&189.15212&62.22800&0.56&11.09&---                &$<$2.06e-17&$<$1.08e-16&---&---&---&---     &---&--- \\\\ 29&188.97445&62.22703&0.88&12.00&---                &$<$6.34e-17&$<$3.07e-16&---&---&---&---     &---&--- \\\\ 30&189.50379&62.22669&0.44&11.77&hdfn-003            &3.72e-14&4.19e-14&43.41&43.46&-0.60$^{+0.01}_{-0.01}$&-0.31$^{+0.01}_{-0.01}$ &BL AGN&AGN\\\\ 31&189.29727&62.22530&2.00&12.82&hdfn-173            &5.75e-17&3.56e-16&42.17&42.96&0.15$^{+0.15}_{-0.14}$&-0.45$^{+0.11}_{-0.13}$ &---     &AGN\\\\ 32&189.37812&62.21623&---&---&---                &$<$4.86e-17&$<$1.43e-16&---&---&---&---     &---&--- \\\\ 33&189.08131&62.21457&0.47&11.52&hdfn-112            &8.15e-17&$<$5.19e-16&40.82&---&-1 &--- &HEL AGN  &AGN\\\\ 34&189.42242&62.21422&1.60&12.52&hdfn-408            &5.57e-17&5.03e-16&41.93&42.88&0.37$^{+0.19}_{-0.18}$&-0.30$^{+0.15}_{-0.21}$&SFG      &AGN\\\\ 35&189.14381&62.21139&1.22&12.56&hdfn-087            &6.15e-17&1.87e-16&41.69&42.17&-0.23$^{+0.15}_{-0.19}$&-0.86$^{+0.15}_{-0.18}$&SFG      &Amb\\\\ 36&189.14571&62.20672&0.56&11.30&---               &$<$3.01e-17&$<$9.47e-17&---&---&---&---     &---&--- \\\\ 37&189.42151&62.20581&---&---&---                &$<$5.30e-17&$<$1.83e-16&---&---&---&---     &---&--- \\\\ 38&189.43598&62.20522&0.23&10.41&---                &$<$1.56e-16&$<$1.65e-16&---&---&---&---     &---&--- \\\\ 39&189.14363&62.20360&0.46&11.65&hdfn-088            &1.47e-16&2.53e-16&41.04&41.28&-0.38$^{+0.10}_{-0.14}$&-2.10$^{+0.12}_{-0.14}$&HEL AGN  &AGN\\\\ 40&189.40440&62.20136&0.41&11.10&hdfn-493            &8.42e-17&$<$3.17e-16&40.69&---&-1&---    &---  &SFG\\\\ 41&189.42067&62.20013&1.17&12.06&hdfn-008            &7.26e-15&1.40e-14&43.71&44.00&-0.39$^{+0.02}_{-0.02}$&0.71$^{+0.02}_{-0.02}$&SFG     &AGN\\\\ 42&189.15318&62.19889&0.56&11.11&hdfn-080            &1.28e-15&2.20e-15&42.18&42.42&-0.48$^{+0.04}_{-0.03}$&-0.20$^{+0.04}_{-0.04}$&SFG      &AGN\\\\ 43&189.27452&62.19824&0.90&11.75&hdfn-413            &$<$4.08e-17&1.18e-17&---&40.65&1&-2.62$^{+0.89}_{-0.83}$&SFG      &Amb\\\\ 44&189.45862&62.19823&---&---&---                &$<$4.75e-17&$<$2.89e-16&---&---&---&---     &--- &---\\\\ 45&189.25661&62.19619&1.76&12.73&hdfn-471            &3.22e-17&$<$1.09e-16&41.79&---&-1&--- &---     &SFG\\\\ \\hline \\hline \\end{tabular} \\end{table} \\addtocounter{table}{-1} \\begin{table} \\centering \\caption{- continued} \\begin{tabular}{|@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l||@{\\hspace{-0.1mm}}l|} \\hline \\hline ID&RA& Dec. & z & $L_{\\rm IR}$ & HDFN ID & $f_{\\rm SX}$   & $f_{\\rm HX}$   & $L_{\\rm SX}$ & $L_{\\rm HX}$ & HR & $f_{\\rm HX}$/$f_{\\rm R}$ &Spectral & X-ray\\\\ & deg & deg &   & log\\,(L$_\\odot$) &         & erg/s/cm$^{2}$ & erg/s/cm$^{2}$ & log (erg/s)     & log (erg/s)     &    &  & Class    & Class\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)&(14)\\\\  \\hline 46&189.03671&62.19546&1.34&12.43&---                &3.86e-17&$<$3.22e-16&41.58&---&-1&---  &---    &SFG\\\\ 47&189.05208&62.19458&0.28&10.79&---                &$<$5.80e-16&$<$3.53e-15&---&---&---&---   &---   &---\\\\ 48&189.22239&62.19434&1.27&12.30&hdfn-056            &9.08e-17&$<$1.66e-16&41.90&---&-1&--- &SFG      &SFG\\\\ 49&189.32612&62.19255&---&---&---                &$<$4.22e-17&$<$3.02e-16&---&---&---&--- &---     &---\\\\ 50&188.88702&62.18744&---&---&---                &$<$7.02e-16&$<$8.41e-16&---&---&---&---&---      &---\\\\ 51&189.01356&62.18636&0.64&11.74&hdfn-299            &1.42e-16&$<$5.94e-16&41.37&---&-0.54$^{+0.24}_{-0.22}$&--- &SFG      &SFG\\\\ 52&188.94873&62.18273&0.45&11.35&---                &$<$2.04e-16&$<$6.08e-16&---&---&---&--- &---     &---\\\\ 53&189.40081&62.18282&---&---&---                &$<$2.49e-16&$<$1.05e-15&---&---&---&--- &---     &---\\\\ 54&189.39812&62.18229&---&---&hdfn-431            &6.44e-17&5.89e-16&---&---&0.25$^{+0.24}_{-0.17}$&-0.57$^{-0.44}_{-0.75}$&---     &AGN\\\\ 55&189.00195&62.18134&---&---&---                &$<$9.20e-17&$<$2.07e-16&---&---&---&--- &---     &---\\\\ 56&189.12131&62.17943&1.01&12.07&hdfn-098            &1.62e-16&1.60e-15&41.91&42.91&0.37$^{+0.07}_{-0.08}$&0.32$^{+0.06}_{-0.06}$&---     &AGN\\\\ 57&188.97516&62.17871&0.51&11.19&---                &$<$6.25e-16&$<$2.90e-15&---&---&---&---  &---    &---\\\\ 58&189.36538&62.17665&0.21&10.50&---                &1.21e-16&$<$3.40e-16&40.20&---&-1&--- &---     &SFG\\\\ 59&189.21302&62.17524&0.41&11.25&hdfn-196            &8.44e-17&$<$2.12e-16&40.69&---&-1&--- &SFG      &SFG\\\\ 60&189.06332&62.16909&1.03&11.82&hdfn-438            &2.65e-17&5.91e-17&41.14&41.49&-0.01$^{0.35}_{-0.27}$&-1.47$^{+0.61}_{-1.41}$&---     &Amb\\\\ 61&189.14032&62.16832&1.02&11.93&hdfn-091            &2.66e-16&3.84e-16&42.13&42.29&-0.46$^{+0.10}_{-0.10}$&-0.78$^{+0.11}_{-0.12}$ &SFG      &AGN\\\\ 62&189.35540&62.16835&0.41&11.06&---                &$<$4.61e-17&$<$1.62e-16&---&---&---&---  &---    &---\\\\ 63&189.13030&62.16607&---&---&hdfn-212            &8.25e-17&5.74e-18&---&---&-0.46$^{+0.25}_{-0.22}$&--- &---    &---\\\\ 64&189.02738&62.16433&0.64&11.66&hdfn-405            &6.54e-18&1.61e-15&40.03&42.42&0.72$^{+0.11}_{-0.09}$&-0.27$^{+0.06}_{-0.07}$&SFG      &AGN\\\\ 65&189.23243&62.15487&0.42&11.17&---                &3.59e-17&$<$2.83e-16&40.34&---&-1&--- &---     &SFG\\\\ 66&189.34521&62.15184&---&---&---                &$<$6.40e-17&$<$1.08e-16&---&---&---&---   &---  &---\\\\ 67&189.29063&62.14482&0.90&12.27&hdfn-480            &2.11e-16&$<$7.28e-16&41.90&---&-1&--- &SFG      &SFG\\\\ 68&189.10174&62.14340&0.95&11.86&---               &6.05e-17&$<$2.81e-16&41.42&---&-1&---   &---   &SFG\\\\ 69&189.13844&62.14298&0.93&11.90&hdfn-093            &2.57e-15&3.21e-15&43.03&43.12&-0.56$^{+0.03}_{-0.03}$&-0.18$^{+0.04}_{-0.04}$&HEL AGN  &AGN\\\\ 70&189.19445&62.14258&0.97&12.11&hdfn-484            &5.52e-17&$<$1.77e-16&41.40&---&-1&--- &SFG      &SFG\\\\ 71&189.23308&62.13562&0.79&11.97&hdfn-328            &8.00e-17&$<$1.16e-16&41.35&---&-1&--- &SFG      &SFG\\\\ 72&189.22781&62.13450&0.84&11.65&---                &$<$2.38e-17&$<$1.06e-16&---&---&---&---     &---&--- \\\\ 73&189.26708&62.13202&1.25&12.24&hdfn-043            &2.73e-15&4.14e-15&43.36&43.54&-0.49$^{+0.03}_{-0.03}$&-0.07$^{+0.03}_{-0.03}$&HEL AGN  &AGN\\\\ 74&189.19095&62.13173&1.43&12.27&hdfn-275            &1.19e-16&$<$4.16e-16&42.14&---&-1&--- &SFG      &SFG\\\\ 75&189.27819&62.12293&---&---&hdfn-181            &6.86e-16&2.78e-15&---&---&-0.05$^{0.01}_{-0.11}$&0.45$^{+0.06}_{-0.06}$&---    &AGN\\\\ 76&189.08724&62.12061&1.15&12.14&---                &$<$9.39e-17&$<$6.36e-16&---&---&---&---     &---&--- \\\\ 77&189.14513&62.11948&0.63&11.35&---                &$<$6.57e-17&$<$1.43e-15&---&---&---&---     &---&--- \\\\ 78&189.15154&62.11859&0.28&10.67&---                &5.44e-17&2.49e-15&40.10&41.76&0.80$^{+0.85}_{-0.72}$&-0.97$^{+0.04}_{-0.04}$&---     &AGN\\\\ 79&189.06666&62.11690&---&---&---                &$<$1.28e-16&$<$6.32e-16&---&---&---&---    &---  &---\\\\ 80&189.13248&62.11211&---&---&---                &5.66e-18&$<$5.69e-16&---&---&-1&---&---      &---\\\\ 81&189.21462&62.11217&0.84&11.72&---                &8.67e-17&$<$3.49e-16&41.45&---&-1&--- &---     &SFG\\\\ 82&189.12056&62.10446&1.26&12.18&hdfn-213            &7.56e-16&7.44e-15&42.81&43.81&0.36$^{+0.07}_{-0.06}$&0.34$^{+0.04}_{-0.04}$&SFG      &AGN\\\\ 83&189.22622&62.09676&0.56&11.53&---                &$<$4.31e-17&$<$2.16e-16&---&---&---&---&---      &---\\\\ 84&189.14024&62.07928&0.52&11.46&---                &$<$7.00e-16&$<$2.29e-15&---&---&---&--- &---     &---\\\\ \\hline \\hline \\end{tabular} \\label{table:all} \\end{table} \\twocolumn"
    },
    "1107/1107.3971_arXiv.txt": {
        "abstract": "In this paper we compare the performances of the $\\chi^{2}$ and median likelihood analysis in the determination of cosmological constraints using type Ia supernovae data. We perform a statistical analysis using the 307 supernovae of the Union 2 compilation of the Supernova Cosmology Project and find that the $\\chi^{2}$ statistical analysis yields tighter cosmological constraints than the median statistic if only supernovae data is taken into account. We also show that when additional measurements from the Cosmic Microwave Background and Baryonic Acoustic Oscillations are considered, the combined cosmological constraints are not strongly dependent on whether one applies the $\\chi^{2}$ statistic or the median statistic to the supernovae data. This indicates that, when complementary information from other cosmological probes is taken into account, the performances of the $\\chi^2$ and median statistics are very similar, demonstrating the robustness of the statistical analysis. ",
        "introduction": "More than a decade ago, type Ia supernovae (SNIa) provided the first clear evidence in favor of cosmic acceleration \\cite{Riess:1998cb,Perlmutter:1998np}. Since then, the availability of ever larger, higher-quality SNIa datasets, as well as measurements using other cosmological probes, such as the Cosmic Microwave Background (CMB) or the Baryonic Acoustic Oscillations (BAO), have been providing overwhelming evidence for the existence of dark energy \\cite{Frieman:2008sn,Amanullah:2010vv,Komatsu:2010fb}, a fluid with large negative pressure capable of driving the acceleration of the universe.  In the so-called $\\Lambda{\\rm CDM}$ model, also known as the concordance model, the dark energy role is played by a cosmological constant $\\Lambda$, responsible for approximately 73\\% of the energy density of the universe at the present day. The remaining percentage is mainly in the form of cold matter, most of which non-baryonic and dark. Radiation is residual at the present time and the universe is spatially flat. Despite the good agreement with observational data, this model has little appeal on theoretical grounds since the value of $\\Lambda$ required to explain the observed cosmic acceleration is off by $\\sim120$ orders of magnitude from the standard quantum field theory prediction. The coincidence between our observing time and the time of the onset of cosmic acceleration is also puzzling \\cite{Avelino:2004vy,Barreira:2011qi}. Many other dark energy models have been proposed, with the most influential being the ones based on dynamical scalar fields. In these models, the energy density varies with time and suitable choices of the scalar field lagrangian can relax some of the problems associated with the cosmological constant (see, for instance, the review \\cite{Copeland:2006wr}). As a result, an important step towards a better understanding of the dark energy involves further testing of its possible dynamical nature.  The determination of dark energy constraints from observational data requires a robust statistical data analysis. In the particular case of SNIa, the usual procedure is to carry out a standard $\\chi^{2}$ likelihood analysis. However, in \\cite{Gott:2000mv,Avelino:2001xj} it has been argued that the median statistic could be a more reliable alternative. The use of the median, despite having the drawback of not being as contraining as the $\\chi^{2}$ analysis, has the strong advantage of requiring weaker assumptions about the data, thus yielding more trustworthy constraints. Moreover, the median is less vulnerable to the presence of {}``outliers'', which is also a significant advantage given current uncertainties about the physics of SNIa.  In this work we revisit and extend the analysis of \\cite{Avelino:2001xj}, updating it using the recent Union 2 SNIa compilation of the Supernova Cosmology Project (SCP) \\cite{Amanullah:2010vv}. We compare the performances of $\\chi^{2}$ and median statistics in the determination of cosmological constraints using SNIa data, considering also CMB and BAO measurements. The layout of this paper is as follows. In Sec. II we describe the application of $\\chi^{2}$ and median statistics to SNIa data. We also discuss the use of additional information contained in the CMB and BAO. In Sec. III we present and discuss our results. Finally, we conclude in Sec. IV. ",
        "conclusions": "In this paper we compared the cosmological constraints derived from the Union 2 SNIa dataset using $\\chi^{2}$ and median statistics. In the absence of CMB and BAO constraints, we have shown that the $\\chi^{2}$ statistic yields tighter cosmological constraints than the median statistic, as a result of the stronger assumptions it makes about the SNIa magnitude error distribution. On the other hand, when CMB and BAO information is taken into account, the performances of both statistics are very similar. Hence, we conclude that the assumption of the gaussian distribution of the errors in the SNIa analysis does not appear to be critical in the determination of cosmological constraints, provided that complementary information from other cosmological probes is also taken into account."
    },
    "1107/1107.2599_arXiv.txt": {
        "abstract": "Theory and simulations suggest that magnetic fields from radio jets and lobes powered by their central super massive black holes can be an important source of magnetic fields in the galaxy clusters. This is paper II in a series of studies where we present self-consistent high-resolution adaptive mesh refinement cosmological magnetohydrodynamic (MHD) simulations that simultaneously follow the formation of a galaxy cluster and evolution of magnetic fields ejected by an active galactic nucleus (AGN). We studied 12 different galaxy clusters with virial masses ranging from 1 $\\times$ 10$^{14}$ to 2 $\\times$ 10$^{15}$ M$_{\\odot}$. In this work we examine the effects of the mass and merger history on the final magnetic properties. We find that the evolution of magnetic fields is qualitatively similar to those of previous studies. In most clusters, the injected magnetic fields can be transported throughout the cluster and be further amplified by the intra-cluster medium (ICM) turbulence during the cluster formation process with hierarchical mergers, while the amplification history and the magnetic field distribution depend on the cluster formation and magnetism history. This can be very different for different clusters.  The total magnetic energies in these clusters are between 4 $\\times$ 10$^{57}$ and $10^{61}$ erg, which is mainly decided by the cluster mass, scaling approximately with the square of the total mass. Dynamically older relaxed clusters usually have more magnetic fields in their ICM. The dynamically very young clusters may be magnetized weakly since there is not enough time for magnetic fields to be amplified. ",
        "introduction": "The increasing detections of radio emission from galaxy clusters, called radio halos and relics \\citep[see][]{Carilli02, Ferrari08, Giovanini09} suggest that the ICM is permeated with magnetic fields. Radio halos are generally diffuse and extended over $\\ge 1$ Mpc, covering the whole clusters, while radio relics, which are often observed at the edges at clusters, can extend as long as 2 Mpc \\citep[e.g.][]{Weeren10}. By assuming that the total energy in relativistic electrons is comparable to the magnetic energy, the magnetic fields in the cluster halos can be estimated at $0.1-1.0$ $\\mu$G and the total magnetic energy can be as high as $10^{61}$ erg \\citep{Feretti99}.  Magnetic fields in galaxy clusters  are also extensively studied through Faraday rotation measurements (FRM). Distribution of FRM, combined with the ICM density measurement, often yields cluster magnetic fields of a few to ten $\\mu$G level, mostly in the cluster centers \\citep{Carilli02}. More interestingly, FRM was investigated to suggest that cluster  magnetic fields may have a Kolmogorov-like turbulent spectrum in core regions \\citep{Vogt05}, with an energy spectrum peak at several kpc. Other studies \\citep[for example][]{Taylor93, Eilek02} have suggested that the coherence scales of magnetic fields can range from a few kpc to a few hundred kpc, implying large amounts of magnetic energy and fluxes in the ICM \\citep{Colgate00}. Recently, study of magnetic fields  by FRM \\citep{Govoni06,Guidetti08,Bonafede10} are extended to the outer part of clusters by observations of more radio galaxies behind or embedded in these clusters. It is expected that the Extened Very Large Array (EVLA) will provide unprecedented new observations on the magnetic fields in the ICM.   It is unlikely that magnetic fields have been dynamically important during the cluster formation. But, it is suggested that the strength and geometry of magnetic fields in clusters may play a crucial role in cluster formation through other processes, such as heat transport, which consequently affect the applicability of clusters as sensitive probes for cosmological parameters \\citep{Voit05}. In addition, since magnetic fields are closely related to synchrotron emission and FRM, the distribution of magnetic fields is important to the understanding the radio observations of the ICM.  Magnetic field evolution is highly nonlinear during cluster formation and difficult  to be studied analytically. Cosmological MHD simulations are being used to study the properties of magnetic fields in the ICM. Such simulations can be very useful in interpreting the magnetic field information from the observation results. Since the origin of cluster magnetic fields is still being debated \\citep{Widrow02},  various initial magnetic fields are used in MHD cluster formation simulations. Simulations were done with initial magnetic fields either from some random or uniform fields at high redshifts \\citep{Dolag02, Dubois08, Dubois09} or from the outflows of normal galaxies \\citep{Donnert08}. All these simulations have found that fields can be further amplified by cluster merger \\citep{Roettiger99} and turbulence in addition to collapse, and their findings roughly match results from observations. On the other hand, very small seed fields from some first principle mechanism, like Biermann battery effect, are also studied \\citep{Kulsrud97,Xu09b} in galaxy cluster simulations. \\citet{Kulsrud97} and \\citet{Ryu08} suggested that very week seed fields can be amplified by dynamo processes in clusters, though current simulations have not been able to model such processes self-consistently.  Magnetic fields in large scale radio jets and lobes from AGNs serve as one very intriguing source of cluster magnetic fields because observations show that they can carry large amounts of magnetic energy and flux \\citep{Burbidge59,Kronberg01,Croston05, McNamara07}.  The magnetization of the ICM and the wider inter-galactic medium (IGM) by AGNs has been suggested on the energetic grounds \\citep{Colgate00, Furlanetto01, Kronberg01}, without details of the physical processes of magnetic field transportation and amplification. Through cosmological MHD simulations, \\citet{Xu09} showed that magnetic fields injected  from a single AGN into a local region can be sufficient to magnetize the whole cluster to the micro Gauss level by the operation of small-scale dynamo \\citep{Brandenburg05,Subramanian06}. Recently, \\citet{Sur10} used MHD simulations to show that same process can also operate during the formation of first stars. Small-scale dynamo may be  an important process to generate magnetic fields in various cosmic objects \\citep{Schleicher10}.  In the first paper \\citep{Xu10} of this study, we studied the magnetic field evolution in a single massive, relaxed galaxy cluster with different AGN injected energies and injection redshifts.  That paper found that, as long as the magnetic fields are injected before the active merger period during the cluster formation history, the AGN magnetic fields can be spread throughout the whole cluster and get substantial amplification. The final magnetic fields are weakly dependent on the amount of initial magnetic fields. But the behavior of magnetic fields in clusters of different masses and at various dynamical states, young (unrelaxed) or old (relaxed) clusters, is still not well studied.  In this paper, we perform a series of high resolution adaptive mesh refinement (AMR) MHD simulations of 12 galaxy clusters of different masses with initial magnetic fields injected by an AGN.  This allows us to investigate the robustness of magnetizing the ICM using the AGN magnetic fields and address additional questions that are not examined with a single cluster. We explore the properties of magnetic field distribution and their evolution in the ICM of those clusters in different dynamical states during their formation histories. The organization of this paper is as follows. In Section \\ref{sec:model}, we provide the details of the simulation setup, including the cluster formation and the magnetic fields injection. We then summarize the main results in Section \\ref{sec:result}. We present the detailed spatial distribution of magnetic fields, evolution of the magnetic energy and the radial profiles of magnetic fields strength. We also present and discuss the properties of synthetic Faraday rotation measurement of our simulated clusters.  In Section 4, we present a summary of the main findings and conclusions. ",
        "conclusions": "In this paper, we report an ensemble of simulations of magnetic field evolution in galaxy clusters with a wide range of masses between 1 $\\times$ 10$^{14}$ and 2 $\\times$ 10$^{15}$ M$_\\odot$ in various dynamical states.  With similar amounts of initial magnetic fields injected from a single AGN at a high redshift, all clusters are eventually filled with micro Gauss magnetic fields, except for dynamically very young clusters. The power spectra of kinetic energy density show that the ICM is in a turbulent state, while the spectra of magnetic energy density show that the small-scale dynamo process is being driven by by the ICM turbulence. This result, along with the previous study of a single cluster with the magnetic field injections of different amounts of magnetic energy and at different redshifts \\citep{Xu10}, suggests that magnetization of galaxy clusters by the operation of small-scale dynamo with the seed magnetic fields from AGN is very robust, and it produces magnetic fields consistent with observed magnetic fields.  The magnetic field evolution and distribution are decided both by the masses of the clusters and their dynamical formation histories. The total magnetic energy is approximately scaled as the square of the virial mass of the cluster, while the dynamically older (relaxed) clusters usually have more magnetic fields. This implies that the cluster magnetic fields are mostly determined by the dynamo process of the ICM turbulence generated by the hierarchical mergers.  Additional magnetic fields from more AGNs or smaller cluster forming halos will not have a major impact on the final magnetic fields in a massive cluster. The $\\gamma$ in the scaling relation between magnetic field and gas density radial profiles ($|B|$ $\\propto$ n$^{\\gamma}$) range between 0.3 and 0.81 for our simulated clusters, while relaxed clusters have flatter magnetic fields profiles. In addition, the relaxed clusters usually have self-similar magnetic field radial distribution, while the field distribution in younger unrelaxed clusters is disturbed by their recently mergers.  Though our simulated clusters only have initial magnetic fields from a single local source, most volumes in the simulated clusters are well magnetized.  We also studied Faraday rotation measurements obtained from the magnetic field and gas density distribution in our simulated clusters. They are also determined by both the cluster sizes and their dynamical states. The radial profiles of $|RM|$ resemble the profiles of the magnetic field strength, and are consistent with the pattern from the observational data. The RM maps reflect the recent cluster mergers, as well as the cluster magnetism history.  There are very long filaments in the RM maps in the recently merged clusters, while their small scale patchy bands reflect the ICM turbulence. Very different distributions of RM of the same cluster but with magnetic fields injected in different locations are observed.  This suggests that RM distribution may be a good probe not only for the cluster formation but also its magnetism history. A detailed study on the relation of RM from simulated clusters with their ICM turbulence will be presented in a forthcoming paper."
    },
    "1107/1107.1740.txt": {
        "abstract": "We present the first systematic study of X-ray flare candidates in short gamma-ray bursts (SGRBs) exploiting the large 6-year \\emph{Swift} database with the aim to constrain the physical nature of such fluctuations. We find that flare candidates appear in different types of SGRB host galaxy environments and show no clear correlation with the X-ray afterglow lifetime; flare candidates are detected both in SGRBs with a bright extended emission in the soft $\\gamma$-rays and in SGRBs which do not show such component. We furthermore show that SGRB X-ray flare candidates only \\emph{partially} share the set of observational properties of long GRB (LGRB) flares. In particular, the main parameter driving the duration evolution of X-ray variability episodes in both classes is found to be the elapsed time from the explosion, with very limited dependence on the different progenitors, environments, central engine life-times, prompt variability time-scales and energy budgets. On the contrary, SGRB flare candidates significantly differ from LGRB flares in terms of peak luminosity, isotropic energy, flare-to-prompt luminosity ratio and relative variability flux. However, these differences disappear when the central engine time-scales and energy budget are accounted for, suggesting that (i) flare candidates and prompt pulses in SGRBs likely have a common origin; (ii) similar dissipation and/or emission mechanisms are responsible for the prompt and flare emission in long and short GRBs, with SGRBs being less energetic albeit faster evolving versions of the long class. Finally, we show that in strict analogy to the SGRB prompt emission, flares candidates fall off the lag-luminosity relation defined by LGRBs, thus strengthening the SGRB flare-prompt pulse connection. ",
        "introduction": "\\label{Sec:Int} With an isotropic peak luminosity up to $10^{54}\\,\\rm{erg\\,s^{-1}}$, gamma-ray bursts (GRBs) are the brightest objects in the $\\gamma$-ray sky during their short lives ($\\Delta t \\sim 0.1-100$ s). Their duration-spectral hardness distribution gives evidence for the presence of two classes (\\citealt{Kouveliotou93}): long  and short GRBs (observed duration longer and shorter than 2 s, respectively), with short bursts appearing slightly harder. The dichotomy in the duration-hardness dimensions suggested separate progenitor populations. However, until a few years ago, the distances, energy and environments of SGRBs (short GRBs) remained highly uncertain due to the poor localisation.  The breakthrough in the study of SGRBs occurred thanks to the rapid slew capabilities of the \\emph{Swift} spacecraft (\\citealt{Gehrels04}) which allowed spectroscopic observations to be performed at very early times. These observations revealed that SGRBs are cosmological, with prompt luminosities comparable to LGRBs albeit significantly less energetic; with similar afterglows (\\citealt{Nysewander09}) but residing in completely different environments. In sharp contrast to LGRBs, short bursts have been localised both in early-type and late-type host galaxies (see \\citealt{Berger11} and references therein), pointing to an \\emph{old} progenitor population. The detection of supernovae associated to LGRBs (see e.g. \\citealt{Kulkarni98}; \\citealt{Stanek03}; \\citealt{Fruchter06} and references therein) provided instead support to models invoking \\emph{young} stellar progenitors. According to the standard scenario, LGRBs originate from the collapse of rapidly-rotating, massive stars (\\citealt{MacFadyen99}), while  SGRBs are believed to result from the coalescence of a binary system of compact objects (neutron star plus neutron star NS+NS or neutron star plus black hole NS+BH, \\citealt{Paczynski86}; \\citealt{Eichler89}; \\citealt{Narayan92}).  Despite fundamental theoretical and observational progress, the nature of SGRB progenitors remains elusive.  Numerical simulations show that the active stage of a NS+NS merger typically lasts $\\sim(0.01-0.1)$ s (see e.g. \\citealt{Nakar07} and references therein\\footnote{In a recent study \\cite{Rezzolla11} found $\\Delta t\\sim 0.3$ s.}): material ejected during the merger is expected to accrete on time-scales of the order of $1-10$ s (the exact value depending on the accreting disk viscosity parameter and details of the ejection process). Thus, the detection of central engine activity on time-scales much longer than the usual dynamical or even viscous time-scales  would challenge the currently accepted scenario (see \\citealt{Nakar07} for a recent review).  Long-lasting ($\\Delta t \\gg 10\\,\\rm{s}$), soft energy tails detected in several SGRBs during their prompt emission  (the so-called extended emission, see \\citealt{Norris10a} and references therein) represent such a case and pose severe constraints to existing models, especially when energetically dominating with respect to the primary burst (\\citealt{Perley09}).  The same is true for the recently discovered presence of precursors (\\citealt{Troja10}). Equally challenging would be the detection of late-time central engine activity in the form of \\emph{flares} superimposed over the smooth SGRB X-ray afterglow.  Flares are currently detected in  $\\sim30$\\% of \\emph{long} GRBs X-ray afterglows (\\citealt{Chincarini10}) as fast-rise exponential-decay features whose spectral and temporal properties have been demonstrated to show a strict analogy to LGRB prompt pulses (\\citealt{Margutti10b}): this finding suggested that flares might originate from re-activations of the LGRB central engine. Several ideas on how to explain the possible presence of flares in \\emph{short} GRBs have been explored as well: the fragmentation of the outer parts of an hyper-accreting disk around the newly formed black hole as a result of gravitational instabilities could potentially lead to large-amplitude variations of the central engine output of both long and short GRBs (\\citealt{Perna06}). Alternatively, the late-time accretion of material launched into eccentric but gravitationally bound orbits during the compact binary merger could provide the fuel to revive the central engine activity (\\citealt{Rosswog07}). The long term evolution of debris following the tidal disruption of compact objects has been identified by \\cite{Lee09} as a feasible mechanism to produce flares. Finally, as an alternative  in the context of accretion-powered models, magnetic halting may also give rise to secondary episodes of delayed activity as suggested by \\cite{Proga06}. However, the observational properties of flares in SGRBs have not been determined, yet, so that it is at the moment unclear if any of these models would be able to explain the observations.  While SGRB X-ray light-curves clearly show temporal variability superimposed over a smooth decay, the presence of real flares in short bursts is questionable. In particular, it is at the moment unclear if what is currently identified as \\emph{short} GRB flare emission (see e.g. \\citealt{LaParola06} for GRB\\,051210) quantitatively shares the very same properties of the population of \\emph{long} GRB flares: are there fast varying $\\Delta t/t\\ll1$, prominent temporal features in the afterglow of SGRBs with properties reminiscent of the long GRB flaring emission? Do SGRB flare \\emph{candidates} follow the entire set of relations found from the analysis of real flares in long bursts? In particular: is the evolution of their temporal and energetic properties compatible with the flare-like behaviour identified by \\cite{Chincarini10}? What is the typical amount of energy released during such episodes of variability? Is there any link between the late-time variability which appears in the X-ray afterglow of SGRBs and their prompt emission? Negligible spectral lag is a defining characteristic of SGRB prompt pulses: is this picture still valid when considering their late-time variability?  Prompted by this set of still open questions, we present the first systematic study of X-ray flare candidates in \\emph{short} GRBs, taking advantage from the large \\emph{Swift} 6-year data-base. Through a homogeneous temporal and spectral analysis of the widest sample of SGRB light-curves available at the time of writing, this study allows us to perform a one-to-one comparison with the properties of X-ray flares detected in long bursts (\\citealt{Chincarini10}, \\citealt{Margutti10b}, \\citealt{Bernardini11}, \\citealt{Margutti11}). The primary goal of this paper is to observationally constrain the origin of SGRB flare candidates providing the reader with a complete picture of their properties.  This work is organised as follows: the sample selection and data reduction is presented in Section \\ref{Sec:DataAn}. Results are described in Section \\ref{Sec:Res} and discussed in Section \\ref{Sec:Disc}. Conclusions are drawn in Section  \\ref{Sec:Con}.  The GRB phenomenology is presented in the observer frame.  Isotropic equivalent luminosities and energies are listed. The observer frame 0.3-10 keV energy band is adopted unless specified. The zero time is assumed to be the trigger time. We use the notation: $Y^{\\rm{SGRB}}_{\\rm{F}}$ ($Y^{\\rm{LGRB}}_{\\rm{P}}$) to indicate that $Y$ refers to the flare (prompt) emission of SGRBs (LGRBs). All the quoted uncertainties are given at 68\\% confidence level (c.l.). Standard cosmological quantities have been adopted: $H_{0}=70\\,\\rm{Km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{\\Lambda }=0.7$, $\\Omega_{\\rm{M}}=0.3$.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{Sec:Con} This work presents the first comprehensive analysis of flare candidates in SGRBs and provides a comparison to the observational properties of X-ray flares in LGRBs with the aim to draw an observational picture of SGRB X-ray variability \\emph{any} theoretical model is required to explain.  Our analysis shows that the origin of the SGRB X-ray light-curve variability is independent from the large-scale host galaxy environment and is not correlated to the detected afterglow life-time. We find that flare candidates appear in different types of SGRB host galaxy environments and show no clear correlation with the X-ray afterglow lifetime; flare candidates are detected both in SGRBs with a bright extended emission (EE) in the soft $\\gamma$-rays and in SGRBs which do not show such component (Sec. \\ref{SubSec:EE}). We cannot exclude that flare candidates originate on top of faint (and undetected in the $\\gamma$-rays) EEs. In particular, SGRB flare candidates are consistent with being drawn from the LGRB flare population when considering: \\begin{enumerate} \\item[1.] The flare to prompt energy ration $E_{\\rm{iso,F}}/E_{\\rm{iso,P}}$ (Fig. \\ref{Fig:Eiso}, Sec. \\ref{SubSec:Energetics}); \\item[2.] The relative variability time scale $\\Delta t/t <1$  (Fig. \\ref{Fig:iokaShort}, Sec. \\ref{SubSec:time}); \\item[3.] The rest-frame flare width evolution with time $w(t_{\\rm{pk}})$ (Fig. \\ref{Fig:wtp}, Sec.  \\ref{SubSec:time}). \\item[4.] The hard-to-soft trend of the emitted radiation  (see e.g. Fig. \\ref{Fig:lc100117A}). \\end{enumerate} The main parameter driving the duration of the episodes of variability is the elapsed time from the explosion, with very limited dependence on the different progenitors, environments, life-times, prompt variability time scales and energy budgets. The origin of the flare $w(t_{\\rm{pk}})$ relation must arise from what is in common for the long and short burst models. From another perspective this result implies that for $t_{\\rm{pk}}>100$ s the flare duration is likely to retain no memory of the variability time-scales of the original prompt mechanism at work. This would explain why the flare to prompt pulse width ratio is different for long and short GRBs.  On the contrary, SGRB flare candidates significantly differ from the standard X-ray flare emission observed in LGRBs  at the same $t_{\\rm{pk}}/(1+z)$ in terms of: \\begin{enumerate} \\item[5.] Peak luminosity $L_{\\rm{pk,F}}^{\\rm{SGRB}}\\sim 0.01 L_{\\rm{pk,F}}^{\\rm{LGRB}}$ (Fig. \\ref{Fig:Ltpk}, Fig. \\ref{Fig:Lpiso}, Sec. \\ref{SubSec:Energetics}); \\item[6.] Isotropic energy $E_{\\rm{iso,F}}^{\\rm{SGRB}}\\sim 0.01 E_{\\rm{iso,F}}^{\\rm{LGRB}}$ (Fig. \\ref{Fig:Eiso}, Sec. \\ref{SubSec:Energetics}); \\item[7.] Flare to prompt luminosity ratio $L_{\\rm{pk,F}}/L_{\\rm{pk,P}}$. Flare candidates in SGRBs are $\\sim100$ times dimmer than in LGRB; \\item[8.] Relative variability flux $\\Delta F/F$ (Fig: \\ref{Fig:iokaShort}): we find $\\Delta F/F\\sim1$ for all SGRB flare candidates (Sec. \\ref{SubSec:flux}); \\item[9.] Lag-luminosity relation: like SGRB prompt pulses, flare candidates show shorter lags than expected from the lag-luminosity relation of LGRB flares (Fig. \\ref{Fig:LagLumFlares}, Sec. \\ref{SubSec:laglum}). \\end{enumerate}  However and more importantly, the differences listed at points 5., 6., 7. and 8. above disappear once the different time scale of evolution of the long and short GRB central engine as well as the different energy scaling of the two systems are properly accounted for  (Fig. \\ref{Fig:Lpktpkrenorm}, Sec. \\ref{SubSubSec:Energetics1} and \\ref{SubSubSec:Energetics2}). This finding provides a connection between the  properties of the detected SGRB X-ray light-curve variability and LGRB flares, suggesting a common, \\emph{internal} origin. As a result, we conclude that similar dissipation and/or emission mechanisms are likely to be responsible for the prompt and flare emission in long and short GRBs, with SGRBs being less energetic albeit faster evolving versions of the long category.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1107/1107.6003_arXiv.txt": {
        "abstract": "We improve estimates of stellar mass and mass-weighted average age of Type Ia supernova (\\snia) host galaxies by combining UV and near-IR photometry with optical photometry in our analysis. Using 206 \\sneia\\ drawn from the full three-year \\sdssii\\ Supernova Survey (median redshift of $z \\approx 0.2$) and multi-wavelength host-galaxy photometry from \\sdss, \\galex, and \\ukidss, we present evidence of a correlation ($1.9 \\sigma$ confidence level) between the residuals of \\sneia\\ about the best-fit Hubble relation and the mass-weighted average age of their host galaxies.  The trend is such that older galaxies host \\sneia\\ that are brighter than average after standard light-curve corrections are made.  We also confirm, at the $3.0 \\sigma$ level, the trend seen by previous studies that more massive galaxies often host brighter \\sneia\\ after light-curve correction. ",
        "introduction": "Observations of Type Ia supernovae (\\sneia) are a key measurement in determining the standard cosmological model.  Their empirical luminosity-distance calibration based on relations between \\snia\\ peak luminosity and both light-curve width and optical colors \\citep{phi93,ham96b,rie96} provides evidence for the accelerated expansion of the Universe and the existence of dark energy \\citep{rie98,per99}. According to the current theory, the progenitor of a \\snia\\ is a carbon-oxygen white dwarf that approaches the Chandrasekhar limit, resulting in a thermonuclear explosion \\citep{whe73,hil00}. However, the exact mechanism by which the progenitor accumulates this mass remains uncertain. Investigations of the physical properties of \\snia\\ host galaxies can provide insight into the environment in which these progenitor systems form. Furthermore, although \\sneia\\ are remarkably standarizable, correcting for light-curve width and color still results in a scatter in peak brightness of $\\sim 0.15$ mag \\citep{guy07,jha07,con08}. Studying how variations in \\snia\\ luminosities depend on the environment of the progenitor will help reveal the origin of this scatter.  Over the years, several correlations between \\sneia\\ and the properties of their progenitors and environments have been discovered.  For example, intrinsically brighter \\sneia\\ tend to occur in galaxies with younger stellar populations while fainter ones often occur in passively evolving galaxies \\citep{ham95,ham00,sul06}. Studies have also shown that per unit stellar mass, the rate of occurrence of \\sneia\\ within a galaxy declines with decreasing SFR \\citep{van90,man05,sul06}. In addition, properties of the progenitors themselves can directly influence light-curve properties of \\sneia. Theoretical models generally agree that the metallicity of the white dwarf progenitor affects the amount of radioactive $^{56}$Ni produced in the thermonuclear explosion, the decay of which powers the light curve of the \\sn\\ \\citep{hof98,tim03}.  Assuming that global metallicity correlates with progenitor metallicity, \\citet{gal05} presented qualitative evidence suggesting that it is more likely that progenitor age, rather than metallicity, is primarily responsible for the variability in \\snia\\ peak luminosity. The true source of this variability has yet to be determined definitively.  More recently, \\citet{gal08} found that early-type host galaxy metallicity is correlated with residuals on the \\sn\\ Hubble diagram around the best-fit cosmology.  The galaxy mass-metallicity relationship \\citep{tre04} has led several authors to investigate whether mass is a proxy for this metallicity trend with Hubble residual (HR). Indeed, the latest studies have shown that more massive galaxies tend to host \\sneia\\ with residuals that are brighter than average after light-curve correction \\citep{kel10,lam10,sul10}. Age is another host property that can be estimated, which might more directly influence \\sn\\ progenitor systems. \\cite{gal08} plotted HR against luminosity-weighted age using optical spectra from 29 early-type host galaxies but found no significant trend.  \\citet{nei09} used optical and UV photometry to calculate luminosity-weighted ages of 166 nearby host galaxies.  They found that for the subsample of 22 low-extinction host galaxies, there was a $2.1 \\sigma$ trend indicating that \\sneia\\ in older hosts have residuals that are brighter than average.  However, when the full sample was used, the trend disappeared.  In this work, we use \\snia\\ host galaxy photometry spanning the ultraviolet, optical, and near-infrared bands, which allows us to constrain stellar masses and ages of host galaxies by comparing the observed photometry to synthetic photometry generated from stellar population synthesis models.  Knowledge of these physical properties of host galaxies can improve our understanding of \\snia\\ progenitors and the diversity of their light curves. ",
        "conclusions": "\\label{Discussion}  We confirm with a significance of $3.0 \\sigma$ the result found by \\citet{kel10}, \\citet{sul10}, and \\citet{lam10} that massive galaxies tend to host overluminous \\sneia.  We also find indication, with a significance of $1.9 \\sigma$, that even after light-curve correction, overluminous \\sneia\\ tend to occur in older stellar populations. We note that the \\citet{nei09} trend of HR with host age was based on luminosity-weighted age while in this work we calculate mass-weighted age, making a direct comparison difficult.  However, the direction of the age trend we see agrees with the \\citet{nei09} trend for their low-extinction hosts. The HR trend with host luminosity-weighted age plotted by \\citet{gal08} is in the opposite direction from the trend we find here, though the significance of their trend is negligible and their methods different. We expect that mass-weighted age is a more unbiased measure of the age of the galaxy because it is not as strongly affected by UV flux from young stars as luminosity-weighted age \\citep[see][Table 1] {lee07}.  Furthermore, the mass-weighted age gives more weight to older stellar populations, to which \\snia\\ progenitors most probably belong, and is therefore more likely to be correlated with \\snia\\ properties than luminosity-weighted age.  The trends we find of HR with mass and age agree with each other in the sense that galaxy mass and age are correlated, with older galaxies generally being more massive.  Based on the mass-age distribution in Figure~\\ref{figAgevMass}, we split our data into two groups: an $\\langle \\mathrm{Age} \\rangle<5$ Gyr group (which encompasses nearly the entire range of masses) and a logMass $>10.2$ group (which encompasses nearly the entire range of ages).  This was done in an effort to investigate the effect of one of the variables (mass or age) on HR while attempting to ``control\" for the other.  In Figure~\\ref{figControl} we plot HR against age for the logMass $>10.2$ group and HR against mass for the $\\langle \\mathrm{Age} \\rangle <5$ Gyr group.   Within these groups we find that the HR correlation with mass is much weaker than in the full sample, and that the HR plot with age is consistent with no correlation. Estimates of host galaxy mass are more robust and have smaller uncertainties than estimates of age\\footnote{Our method yields uncertainties of $2\\%$ for mass and 27\\% for age, though we emphasize that these uncertainties are \\emph{statistical} only (as described in Section \\ref{DerGalProps} and discussed in Section \\ref{CoaddVsingle}) and that systematic uncertainties on mass are around 0.1 dex (25\\%) at best.  Running our linear regression on the HR vs. mass plot with mass uncertainties inflated to be at least 0.1 dex does not change our results.}, and this in part may be why HRs are more strongly correlated with mass.  Thus, the true underlying property influencing the \\snia\\ explosion is still unclear. The strength of both correlations may be improved by having UV and near-IR matches for all \\sdss\\ hosts and by calculating galaxy magnitudes in a consistent manner through matched apertures for all survey types (as was done, e.g., in \\citealp{hil10}), ensuring that UV-optical-NIR colors are accurate. It is also possible that the relationship between HR and age or mass is not simply linear and may be more complex. Additionally, in an ideal case of a well-resolved extended host, a local age computed from photometry obtained from the location of the \\sn\\ in the galaxy would be preferable to the global galaxy average age that we compute here and would likely correlate more strongly with properties of the \\snia.  \\begin{figure*}[tbp] \\epsscale{1.15} \\plottwo{f7a.eps}{f7b.eps} \\caption{\\emph{Left}: HR versus host galaxy age for the logMass $>10.2$ group. \\emph{Right}:  HR versus host galaxy mass for the $\\langle \\mathrm{Age} \\rangle <5$ Gyr group. The squares are the binned averages.} \\label{figControl} \\end{figure*}   We find that the HR trend with mass persists if we consider the subsample for which \\sdss\\ is complete. Therefore, based on our measurements and the completeness of our dataset, it is likely that \\sneia\\ that are underluminous after light-curve correction do not occur in massive galaxies. This view is consistent with the results of \\citet{kas09} who showed using \\snia\\ simulations and the \\citet{tim03} model that metallicity can affect the explosion physics in such a way as to cause metal-rich progenitors to produce \\sneia\\ that are fast-declining and intrinsically fainter at peak. The usual light-curve correction technique does not account for this metallicity effect, and so metal-rich progenitors result in overluminous \\sneia\\ after corrections for light-curve shape.  Thus the \\citet{kas09} result is in agreement with the trends we see with mass (and, for the full sample, age) since galaxies with higher metallicity are, in general, older and more massive \\citep{tre04,glz05}.  Given our dataset and the large scatter in our trends of HR with age and mass, it is not possible to determine for certain what host property most influences \\snia\\ luminosity.  It is improbable that host mass itself, though better estimated, has a direct impact on \\sneia.  Rather, it is more likely that host mass is correlated with other properties of the host that do directly influence the progenitors of \\sneia. The complexity of the relationships between galaxy properties such as age, mass, metallicity, dust, and SFR makes disentangling the factors that affect \\sneia\\ a challenge. Further study is needed to truly ascertain the origin of these correlations between host properties and Hubble residuals and to potentially pinpoint the cause of these observed trends."
    },
    "1107/1107.1286_arXiv.txt": {
        "abstract": "We present a model of the Galactic Habitable Zone (GHZ), described in terms of the spatial and temporal dimensions of the Galaxy that may favour the development of complex life.  The Milky Way galaxy is modelled using a computational approach by populating stars and their planetary systems on an individual basis using Monte-Carlo methods.  We begin with well-established properties of the disk of the Milky Way, such as the stellar number density distribution, the initial mass function, the star formation history, and the metallicity gradient as a function of radial position and time.  We vary some of these properties, creating four models to test the sensitivity of our assumptions. To assess habitability on the Galactic scale, we model supernova rates, planet formation, and the time required for complex life to evolve.  Our study improves on other literature on the GHZ by populating stars on an individual basis and by modelling SNII and SNIa sterilizations by selecting their progenitors from within this preexisting stellar population.  Furthermore, we consider habitability on tidally locked and non-tidally locked planets separately, and study habitability as a function of height above and below the Galactic midplane.  In the model that most accurately reproduces the properties of the Galaxy, the results indicate that an individual SNIa is $\\sim$5.6$\\times$ more lethal than an individual SNII on average.  In addition, we predict that $\\sim$1.2\\% of all stars host a planet that may have been capable of supporting complex life at some point in the history of the Galaxy.  Of those stars with a habitable planet, $\\sim$75\\% of planets are predicted to be in a tidally locked configuration with their host star.  The majority of these planets that may support complex life are found towards the inner Galaxy, distributed within, and significantly above and below, the Galactic midplane. ",
        "introduction": "\\par Studies of habitability on the Galactic scale are gaining attention as extrasolar planetary searches and advanced models of the evolution of the Milky Way galaxy improve our understanding of the prerequisites for life.  The GHZ is the area in the Galaxy that may favour the emergence of complex life.  The inner boundary of the GHZ is thought to be defined by hazards to planetary biospheres, and the outer boundary is set by the minimum amount of metallicity required for planet formation \\citep{2001Icar..152..185G}.  Given these factors, the GHZ is constrained by the morphology, stellar populations, and chemical evolution of the Milky Way.  The GHZ is defined as a concept analogous to the Circumstellar Habitable Zone (HZ), which is the region around a star where water can remain in a liquid state on a planet's surface.  \\par To assess habitability on the Galactic scale, we model planet formation, supernova rates, and the time required for complex life to evolve.  Estimates in predicting the number of habitable planets in the Galaxy are in their infancy at present.  Candidate habitable planets have been identified in the Gliese 581 system \\citep{2007A&A...476.1373S,2010ApJ...723..954V}, and in the Kepler list of candidate exoplanets \\citep{2011arXiv1102.0541B}. Models can provide insight into the number and distribution of habitable planets that are likely to exist in the Galaxy.  Prior studies suggest that there is a correlation between high-metallicity environments and planet formation \\citep{2004ApJ...616..567I,2004A&A...415.1153S,2005ApJ...622.1102F,2007ApJ...669.1220G,2009ApJ...697..544S,2010PASP..122..905J}.  The metallicity gradient found in the disk of the Galaxy has been measured, and suggests that the number of planets depends on radial position.  Regions greatly affected by the detrimental impacts of supernovae (SNe) are defined by proximity to these transient events.  Therefore, the stellar density and star formation history (SFH) directly influence the frequency at which a planet's biosphere is negatively affected.  Finally, in combination with the periods that SNe are expected to sterilize life in surrounding solar systems, enough time must elapse such that complex life can evolve. \\par Research on the GHZ by \\citet{lineweaver-GHZ} suggests that the GHZ is an annular region that expands with time.  However, \\citet{2008SSRv..135..313P} argues that the entire disk of the Galaxy may be suitable for complex life.  From our work we predict that at the present time, the entire disk of the Galaxy is hospitable to complex life; however, the regions that permit the majority of complex life are contrary to other predictions in the field. \\par Predicting where complex life might be favoured in the Galaxy follows from making assumptions about the prerequisites for life.  Whether or not life exists beyond the context of our solar system cannot be validated at present.  Our definitions of habitability in the Galaxy do not imply any claims of inhabitance. \\par We model the GHZ with the factors described above.  We begin in \\S~\\ref{model} by outlining the construction of our model Milky Way galaxy.  In \\S~\\ref{model_habitability_MWG} we describe the factors that influence habitability on the Galactic scale.  These factors include: the dangers caused by SN sterilizations, how the metallicity gradient influences planet formation, and the timescale requirements for the emergence of biologically complex life.  Results and a comparison with other studies are discussed in~\\S~\\ref{results}.  Finally, in~\\S~\\ref{conclusion} we present the conclusions drawn from our study of habitability in the Galaxy. ",
        "conclusions": "\\par We present a model of habitability within the Milky Way to predict the region(s) in the Galaxy that are expected to favour the emergence of complex life.  We do not find that the inner Galaxy is entirely inhospitable to life; in fact, the greatest number of habitable planets are found in this region.  The metallicity in the inner Galaxy produces a high planet formation rate for long timescales that dominates the negative impact SNe have on habitability.  Furthermore, we observe that over all epochs in our favoured Model (Model~4), 1.2\\% of all stars in the Galaxy host a habitable planet (including both tidally locked and non-locked configurations).  The fraction of stars with a habitable planet ranges from $\\sim$0.25\\% in the outer Galaxy to $\\sim$2.7\\% in the inner Galaxy.  Considering that we find the greatest number of habitable planets to be located in and around the midplane at R$\\sim$2.5 kpc, and that the greatest fraction of stars that support habitable planets exist above and below the midplane at this radial position, we find that the inner region of the Galaxy may support the greatest number of planets conducive to complex life.  More specifically, our findings are at odds with the notion that the GHZ is shaped like an annular ring at R$\\sim$7-9 kpc at the present day, as found by \\citet{lineweaver-GHZ}.  Similarly to~\\citet{2008SSRv..135..313P}, we find the greatest number of habitable planets to exist in the inner Galaxy, with the exception that we observe habitable planets at this radial position to be strongly dependent on height above and below the Galactic midplane. \\par Given that we find the greatest concentration of habitable planets to be located in the inner Galaxy, further research on the disk and overlapping bulge at R$\\lesssim$2.5 kpc is warranted.  The inner boundary of the GHZ is defined by hazards to a planet's biosphere \\citep{2001Icar..152..185G}; however, we find that an inner boundary does not exist when modelling the Galactic disk at R$\\gtrsim$2.5 kpc.  By modelling the region at R$\\lesssim$2.5 kpc, we might find an inner boundary, or we might not find a boundary at all. \\par \\par \\par Studies of habitability on the Galactic scale will improve in the future as Earth-like planets are detected from studies such as the Kepler mission \\citep{2008IAUS..249...17B}, and will benefit from an improved understanding of our Galaxy from the GAIA mission \\citep{2001A&A...369..339P}, and others. \\par The results of planet finding missions will yield estimates of the total number of planets in the HZ of their host stars across the Galactic disk.  The total number of habitable planets estimated here is not the focus of our research; rather, we highlight and predict that these planets exist preferentially in the inner Galaxy, and that they have the ability to survive SN sterilization events for periods conducive to the rise of biologically complex life."
    },
    "1107/1107.3884_arXiv.txt": {
        "abstract": "The $K_s$ band differential star count of the Two Micron All Sky Survey (2MASS) is used to derive the global structure parameters of the smooth components of the Milky Way. To avoid complication introduced by other fine structures and significant extinction near and at the Galactic plane, we only consider Galactic latitude $|b|>30^{\\circ}$ data. The star count data is fitted with a three-component model: double exponential thin disk and thick disk, and a power law decay oblate halo. Using maximum likelihood the best-fit local density of the thin disk is $n_0=0.030\\pm 0.002$ stars/pc$^3$. The best-fit scale-height and length of the thin disk are $H_{z1}=360\\pm 10$ pc and $H_{r1}=3.7\\pm 1.0$ kpc, and those of the thick disk are $H_{z2}=1020\\pm 30$ pc and $H_{r2}=5.0\\pm 1.0$ kpc, the local thick-to-thin disk density ratio is $f_2=7\\pm 1$\\%. The best-fit axis ratio, power law index and local density ratio of the oblate halo are $\\kappa=0.55\\pm 0.15$, $p=2.6\\pm 0.6$ and $f_h=0.20\\pm 0.10$\\%, respectively. Moreover, we find some degeneracy among the key parameters (e.g., $n_0, H_{z1}, f_2$ and $H_{z2}$). Any pair of these parameters are anti-correlated to each other. The 2MASS data can be well-fitted by several possible combinations of these parameters. This is probably the reason that there is a wide range of values for the structure parameters in literature similar to this study. Since only medium and high Galactic latitude data are analyzed, the fitting is insensitive to the scale-lengths of the disks. ",
        "introduction": "In the eighteenth century, the famous astronomer William Herschel showed us the powerful method of star count to understand our own Milky Way \\citep{Hers1785}. The technique has been used since then by generations of astronomers. With great improvement on data collection over the years, more and more details of our Milky Way were unfolded. However, the characteristic scales of smooth Galactic structures (i.e., disks and halo) obtained by previous studies does not converge to a common value as an outcome of improving data collection (see Table~\\ref{table_ref}). The spread of values is attributed to the degeneracy of Galactic model parameters (i.e., same star count data could be fitted equally well by different Galactic models) \\citep{Chen2001, Siegel2002, Juric2008, Bilir2008}. This is due to the different sky regions and limiting magnitudes (i.e., limiting volumes) used in these studies \\citep{Siegel2002, Karaali2004, Bilir2006a, Bilir2006b, Juric2008}. On the contrary, \\citet{Bilir2008, Yaz2010} did not show such degeneracy in determining the Galactic model parameters. This controversy is still actively debated. Therefore, systematic all sky surveys with deeper limiting magnitude and wider sky region, such as the Two Micron All Sky Survey \\citep[2MASS;][]{Skrutskie2006}, the Sloan Digital Sky Survey \\citep[SDSS;][]{York2000}, the Panoramic Survey Telescope \\& Rapid Response System \\citep[Pan-Starrs;][]{Kaiser2002} and the GAIA mission \\citep{Perryman2001} could provide a good opportunity for us to study our Galaxy from a global perspective. Free from limited sky fields, astronomers can acquire many more information from the stellar distribution of these surveys.  On Galactic structure study, besides the simple and smooth two-component model \\citep{Bahcall1980} or the three-component model \\citep{Gilmore1985}, many more structures have been discovered, such as inner bars in the Galactic center \\citep{Alves2000, Hammersley2000, vanLoon2003, Nishiyama2005, Cabrera2008}, and flares and warps \\citep{Lopez2002, Robin2003, Momany2006, Reyle2009}, which has been contributed the variation of disk model parameters with Galactic longitude \\citep{Cabrera2007, Bilir2008, Yaz2010}. Moreover, the overdensities in the halo, such as Sagittarius \\citep{Majewski2003}, Triangulum-Andromeda \\citep{Rocha2004, Majewski2004}, Virgo \\citep{Juric2008}, and in the outer disk, such as Canis Major \\citep{Martin2004}, Monoceros \\citep{Newberg2002,Rocha2003}, show the complexity of the Milky Way. The formation history of our Galaxy is more complicated than what we thought previously.  On stellar luminosity function study, a recent study by \\citet{Chang2010} using the $K_s$ band star count of 2MASS point source catalog \\citep[2MASS PSC;][]{Cutri2003} verified for the first time the universality hypothesis of the luminosity function (i.e., the common practice of assuming one luminosity function for the entire Milky Way).  We are interested in the global smooth structure of our Galaxy. In view of the existing fine structures (e.g., flares, warps, overdensities\\dots etc.), the structure of the smooth components can be determined either by including these fine structures in a grand Galaxy model or by avoiding the sky area ``contaminated'' with these features. The first method demands a complex model, which involves many more structure parameters, and needs high computing power to accomplish the fitting task. \\citet{Lopez2002} is a good example of this method using 2MASS data. The second method is clearly simpler but needs justifications. We observe that (1) the fine structures (e.g., inner bars, flares and warps), which have observable contribution on 2MASS star count data, are all confined in the Galactic plane region. (2) The overdensities or substructures in the outer disk region and the halo are difficult to identify in general. Their contribution is negligible in the 2MASS star count data. Here we quote \\citet{Majewski2004} on halo substructure: ``This substructure is typically subtle and obscured by a substantial foreground veil of disk stars, eliciting its presence requires strategies that optimize the substructure signal compared to the foreground noise.''  We prefer the second method in this paper and only use Galactic latitude $|b|>30^{\\circ}$ 2MASS $K_s$ band star count data to obtain the structure parameters of a three-component model. In addition, the influence of the near infrared extinction in these regions is small. This also allows us to use a simpler extinction model for correction. We describe our model in section 2 and the analysis method in section 3. Section 4 provides the results and a discussion. ",
        "conclusions": "Table~\\ref{result} lists our best-fit results with the corresponding uncertainties. Fig.~\\ref{2Dcon} shows contour plots of likelihood against different pairs of parameters. The contour changes dramatically along the key parameters $n_0$, $H_{z1}$, $f_2$ and $H_{z2}$, but relatively mild along other parameters $H_{r1}$, $H_{r2}$, $f_h$, $p$ and $\\kappa$. This indicates the importance of the key parameters in determining the best-fit result.  Degeneracies exist between some pairs of key parameters, such as $(n_0, H_{z1})$, $(n_0, f_2)$, $(H_{z1}, f_2)$\\dots etc. Here degeneracy means that the likelihood value stays almost the same when the pairs of parameters change together in a particular way. Similar degeneracy between the local thick-to-thin disk ratio $f_2$ and the scale-height of the thick disk $H_{z2}$ has been reported in \\citet{Chen2001,Siegel2002,Juric2008,Bilir2008}. Consequently, it is possible that different combinations of parameters can be regarded as `acceptable' fitting. For example, if we choose a higher local stellar density $n_0$, then we can pick a smaller $H_{z1}$ such that the likelihood value is very close to the maximum likelihood value and assign it as the `best-match' scale-height of the thin disk. Therefore, a thin light color diagonal strip shows on the $n_0$ against $H_{z1}$ contour plot in Fig.~\\ref{2Dcon}. Besides, similar trends happen in other pairs of key parameters. When one parameter of the pair is higher, we can get a similar likelihood value by lowering the other parameter (see the corresponding two-parameter contour plots of key parameters in Fig.~\\ref{2Dcon}). This anti-correlation is not unexpected. The number of stars along the line of sight in the model increases when any one of the key parameters increases. Thus for a given observed number of stars, an increase in one key parameter can be compensated by a decrease in the other. Perhaps this is the reason that our best-fit scale-height of the thin disk, $H_{z1}=360$ pc, is somewhat larger than the reported values, $H_{z1}=285$ pc, in \\citet{Lopez2002} study (they also use star count of 2MASS to obtain Galactic model parameters). Our best-fit local stellar density from $K_s=$-8 to 6.5 mag is $n_0\\sim 0.030$ star/pc$^3$, which is about half of $\\sim 0.056$ star/pc$^3$ of the corresponding value cited in \\citet{Eaton1984}, and a bit lower than $\\sim 0.032$ star/pc$^3$ of the corresponding value cited in \\citet{Lopez2002}. As a result, our fitting tends to choose a larger scale-height of the thin disk. If we force our local stellar density to be comparable with that of \\citet{Lopez2002}, then the corresponding `best-fit' scale-height of the thin disk would be $\\sim$320 pc, which is closer to their result. If we choose even higher local stellar density, then the `best-fit' scale-height of the thin disk would be made between 200 to 300 pc, which is similar to the most of recent studies (see Table~\\ref{table_ref}). Moreover, we do not apply binarism correction in our analysis, and it has been shown that scale-length and scale-height might be underestimated without binarism correction \\citep{Siegel2002, Juric2008, Ivezic2008, Yaz2010}.   For our purpose, we deem that in order to lift the degeneracy it is crucial to have a reliable near infrared luminosity function by observation or a near infrared local stellar density. Unfortunately, a systematic study in this direction is yet to come. Some related studies, such as synthetic luminosity function \\citep[see e.g.,][]{Girardi2005} or luminosity function transformed from optical observation \\citep[see e.g.,][]{Wainscoat1992} do exist, but some uncertainties still need to be settled (e.g., the initial mass function, mass-luminosity relation for NIR and color transformation between different wavelengths\\dots etc.). Once the `true' local stellar density is known, the `true' structure of the Milky Way would be revealed.    In order to see how good the agreement between 2MASS data and our best-fit model, we show an all sky map of the ratios of observed to predicted integrated star count from $K_s=$5 to 14 mag for each node as a function of position on the sky in Fig.~\\ref{comparison}, and the color indicates the values of the ratios. We do not see obvious deviation in the Galactic latitude $|b|>30^{\\circ}$ areas, but only in the Large Magellanic Cloud and the Small Magellanic Cloud areas. For comfirmation, we plot integrated star count from $K_s=$5 to 14 mag along Galactic longitude and latitude for $|b|>30^{\\circ}$ areas in Figs.~\\ref{surf_b} \\& \\ref{surf_l}. We see 2MASS data and our best-fit model agree well and only some small deviations in the nodes at the anti-Galactic center $b\\sim30^{\\circ}$ areas. The significant spikes in Figs.~\\ref{surf_b} \\& \\ref{surf_l} are due to the populations of the Large Magellanic Cloud and the Small Magellanic Cloud. For testing how these small deviations affect the key parameters selection, we exclude these small deviations and re-analyze the data with fixed non-key parameters. It shows only the scale-height of thick disk shifts slightly but still within our error estimation. Besides, we do not see any significant differences of reported overdensities in sky regions with $|b|>30^{\\circ}$. Thus, we conclude that the three-component model can describe the Milky Way structure sufficiently well for high Galactic latitude regions and the single power law luminosity function of \\citet{Chang2010} is a good approximation as well. Since our main purpose is to search for the global Galactic model parameters, we do not try to explore the best-fit result for each node individually as what \\citet{Cabrera2007, Bilir2008, Yaz2010} have done in their studies to seek the variations of disk model parameters with Galactic longitude. Instead, we treat the differences of 2MASS data to our best-fit model as deviations from a global smooth distribution. We believe that similar variations in disk model parameters will be obtained if we take similar analysis procedure (i.e., searching best-fit parameters for each node), but this is beyond the scope of this work. Because we do not consider flares, warps and other overdensities in our model, there are some discrepancies between 2MASS data and the model in the low Galactic latitude regions (see Fig.~\\ref{comparison}). These fine structures make the star distribution more fluffy in the vertical direction toward the edge of the Milky Way. Hence the discrepancy between 2MASS data and the model increases vertically towards the anti-Galactic center region. In addition, the absence of Galactic bulge in our model contributes to the large discrepancy in Galactic center areas. The difference in the low Galactic latitude region needs more delicate analysis to rectify \\citep[see, e.g.,][]{Lopez2002, Momany2006}.  \\subsection{Summary} In summary, we set forth to study the global smooth structure of the Milky Way by a three-component stellar distribution model which comprises two double exponential disks (one thin and one thick) and an oblate halo. The $K_s$ band 2MASS star count is used to determine the structure parameters. To avoid the complication introduced by the fine structures and complex extinction correction close to Galactic plane, we use only Galactic latitude $|b|>30^{\\circ}$ data. There are 10 parameters in the model, but only four of them play the dominant  role in the fitting process. They are the local stellar density of the thin disk $n_0$, the local density ratio of thick-to-thin disk $f_2$, and the scale-height of the thin and thick disks $H_{z1}$ and $H_{z2}$. The best-fit result is listed in Table~\\ref{result}. In short the scale-height of the thin and thick disks are $360\\pm 10$ pc and $1020\\pm 30$ pc, respectively; the scale-length of the thin disk is $3.7\\pm 1.0$ kpc and that of thick disk is $5.0\\pm 1.0$ kpc (the uncertainty in scale-length is large because it is not very sensitive to high latitude data.) The local stellar density ratio of thick-to-thin disk and halo-to-thin disk are $7\\pm 1$\\% and $0.20\\pm 0.10$\\%, respectively. The local stellar density of the thin disk is $0.030\\pm 0.002$ stars/pc$^3$.  An all sky comparison of the 2MASS data to our best-fit model is shown in Fig.~\\ref{comparison}. A good agreement in the Galactic latitude $|b|>30^{\\circ}$ areas is expected from our fitting procedure. In low Galactic latitude regions, fine structures (such as flares, warps\\dots etc.) increase the effective scale-height towards the edge of the Milky way. This is reflected in the fan-like increase in discrepancy towards the anti-Galactic center regions.  Degeneracy (i.e., different combinations of parameters give similar likelihood values) is found in pairs of key parameters (see Fig.~\\ref{2Dcon}). Thus different combinations of parameters may fit the data almost as good as the best-fit one, and these are all legitimate `acceptable' fitting in view of the uncertainty. Therefore, accompanying our lower local stellar density $0.030$ stars/pc$^3$ (from $K_s=-$8 to 6.5 mag) is a higher thin disk scale-height $360$ pc. In the context of NIR star count, the NIR luminosity function or the NIR local stellar density is imperative to determine the scale-height and other Milky Way structure parameters. We hope that systematic study on the luminosity function and the local stellar density in near infrared will be available in the near future."
    },
    "1107/1107.3619.txt": {
        "abstract": "We present 4 model series of the {\\small CODEX} dynamical opacity-sampling models of Mira variables with solar abundances, designed to have parameters similar to $o$~Cet, R~Leo and R~Cas. We demonstrate that the {\\small CODEX} models provide a clear physical basis for the molecular shell scenario used to explain interferometric observations of Mira variables. We show that these models generally provide a good match to photometry and interferometry at wavelengths between the near-infrared and the radio, and make the model outputs publicly available. These model also demonstrate that, in order to match visible and infrared observations, the Fe-poor silicate grains that form within 3 continuum radii must have small grain radii and therefore can not drive the winds from O-rich Mira variables. ",
        "introduction": "Asymptotic Giant Branch (AGB) stars represent the final fusion-powered stage in the evolution of solar-type stars, and the engine by which the vast majority of the material in our Galaxy is recycled from stars back to the interstellar medium \\citep{Gail03}. Mira variables represent the final stage in AGB evolution before they become dust-enshrouded and difficult to observe. They are so bright in the infrared that they can be used as competitive extragalactic distance indicators and probes of star formation history \\citep{Rejkuba04,Menzies10}. They are also unique amongst stellar classes in their opportunity for detailed observations: light curves that differ in shape and amplitude in different bandpasses, photospheres that show different sizes and structure as a function of wavelength, velocity-resolved motions and complex spectra.  The observational literature on Mira variables is very extensive, as partly detailed below in Section~\\ref{sectInterferometry}. In order to make sense of these observations, comprehensive models are required that link physical parameters to pulsation, and pulsation to observed properties and mass-loss. The previous generation of models \\citep[e.g.][]{Hofner98,Hofmann98} suffered from grey or mean-opacity like approximations in their radiative transfer codes, so were not ideally suite to interpreting many observed properties, for example visible brightness or high-resolution spectroscopy. The next generation of models (Upsalla:\\citealt{Hofner03},{\\small CODEX}:\\citealt{Ireland08}) are better suited for modern observational comparisons, but extensive grids have not yet been produced, both because of solvable but difficult computational issues and the lack of a clear calibration for model parameters.  Here we present four physical model series for M-type Mira variables, as a first step in tuning and testing the {\\small CODEX} models to derive physical parameters of Mira variables from observations, and to gain physical insight into the dominant physical processes in Mira-variable atmospheres. Observational predictions including predictions for infrared interferometry of the model series are made available online so that new observations can easily be compared to these models.  One of these model series has parameters based on $o$~Cet, one based on $R$~Leo and two are based on R~Cas with different assumptions. For all models, only the pulsation period is guaranteed to match observations, and in this paper we aim to examine the other model outputs in order to more closely target model series to real stars in future papers. The model construction was described in Paper~I \\citep{Ireland08}: they begin with input parameters of mass, luminosity and composition, with three other free parameters being microturbulent velocity, mixing length ($\\alpha_m$) and turbulent viscosity ($\\alpha_\\nu$). Pulsation is self-excited (i.e. it occurs spontaneously), and the temperature of all layers is calculated by solving the conservation of energy equation via a detailed opacity-sampling method.  Details of the parameter choices for the model series are given in Section~\\ref{sectMethods}, as well as basic comparison of the model light curves to observations. Model predictions for spectra, in particular the effect of extension on spectra, are tested in Section~\\ref{sectSpectra}, and the models are compared to observations of $o$~Cet including infrared and radio interferometry in Section~\\ref{sectInterferometry}.  In Section~\\ref{sectShell}, we compare the model structures to previously published ad-hoc molecular shell papers, and in Section~\\ref{sectRadAccel} we discuss the mass loss rates of the models and the driving mechanisms. In Section~\\ref{sectParams} we discuss the effect of input parameters on the models, and the possibility for better calibrating the input parameters so that, e.g. the mass of individual Miras could be inferred from models. Finally, in Section~\\ref{sectConclusions} we conclude and discuss plans for future work. ",
        "conclusions": "\\label{sectConclusions}  The atmospheric models presented here, based on self-excited pulsation models and opacity-sampling treatment of radiation transport, provide a fairly realistic approximation of the atmospheric density-temperature stratification. Many spectral features are predicted with satisfactory accuracy but some, like TiO bands, require further improvement of the models e.g. the non-LTE treatment of Paper~I. The present sequence of models also comprises only 4 combinations of basic stellar parameters at only a single composition and, therefore, can only describe a relatively small subset of Mira variables. Predictions of the 4 model series presented here (Table 1) are available on-line (Section~\\ref{sectUsingModels}).  There are several potential causes of substantial deviation from spherical asymmetry in Mira variables - including convective cells, weak chaos and Rayleigh-Taylor instabilities \\citet{Woitke06}. Asymmetries are clearly not considered in our model series as they are spherically symmetric, but asymmetries are relatively common in Mira variables when observed at sufficient angular resolution (e.g. \\citet{Ragland06}). The one prediction that can be made from the models is that the wavelengths most susceptible to cycle-to-cycle variations (e.g. L-band where the water shells are optically-thick) should also show asymmetries, as high layers on opposite sides of the star should not be strongly causally connected and show weak chaos. We finally suggest that an observation so far missing in the literature is the astrometric motions of the radio photosphere over several cycles, which should be a strong indicator for the degree of high-layer chaos in Mira atmospheres.  Determination of the internal fundamental model parameters, i.e. mixing-length $\\alpha_m$ and turbulent viscosity $\\alpha_\\nu$ (Table 1), would require observation of a set of stars with different mass, luminosity and pulsation period. For a given pulsation period, a higher-mass star must have a higher luminosity (radius), but a similar effective temperature. We suggest that further studies of low-amplitude pulsators with Mira-like periods such as R~Dor and W~Hya may provide the key to tuning the mixing length parameter $\\alpha_m$ of Mira model series. These are likely higher-mass stars, but with a similar effective temperatures to $o$~Cet. Tuning the turbulent viscosity parameter will best be done by fitting to amplitudes of models with the best-known masses such as, e.g., those kinematically associated with the thick disk."
    },
    "1107/1107.3556_arXiv.txt": {
        "abstract": "{Recently, a new framework for describing the multiverse has been proposed which is based on the principles of quantum mechanics. The framework allows for well-defined predictions, both regarding global properties of the universe and outcomes of particular experiments, according to a single probability formula.  This provides complete unification of the eternally inflating multiverse and many worlds in quantum mechanics.  In this paper we elucidate how cosmological parameters can be calculated in this framework, and study the probability distribution for the value of the cosmological constant.  We consider both positive and negative values, and find that the observed value is consistent with the calculated distribution at an order of magnitude level.  In particular, in contrast to the case of earlier measure proposals, our framework prefers a positive cosmological constant over a negative one.  These results depend only moderately on how we model galaxy formation and life evolution therein.}  \\end{center} \\end{titlepage} ",
        "introduction": "\\label{sec:intro}  An explanation of a small but nonzero cosmological constant is one of the major successes of the picture that our universe is one of the many different universes in which low energy physical laws take different forms~\\cite{Weinberg:1987dv}.  Such a picture is also suggested theoretically by eternal inflation~\\cite{Guth:1982pn} and the string landscape~\\cite{Bousso:2000xa}.  This elegant picture, however, has been suffering from the predictivity crisis caused by an infinite number of events occurring in eternally inflating spacetime.  To make physical predictions, we need to deal with these infinities and define physically sensible probabilities~\\cite{Guth:2000ka}.  Recently, a well-defined framework to describe the eternally inflating multiverse has been proposed based on the principles of quantum mechanics~\\cite{Nomura:2011dt}.  In this framework, the multiverse is described as quantum branching processes viewed from a single ``observer'' (geodesic), and the probabilities are given by a simple Born-like rule applied to the quantum state describing the entire multiverse.  Any physical questions---either regarding global properties of the universe or outcomes of particular experiments---can be answered by using this single probability formula, providing {\\it complete unification} of the eternally inflating multiverse and many worlds in quantum mechanics. Moreover, the state describing the multiverse is defined on the ``observer's'' past light cones bounded by (stretched) apparent horizons; namely, consistent description of the {\\it entire} multiverse is obtained in these limited spacetime regions.  This leads to a dramatic change of views on spacetime and gravity.  In this paper we present a calculation of the probability distribution of the cosmological constant in this new framework of the quantum multiverse.% \\footnote{For earlier studies of the cosmological constant in the context of geometric cutoff measures, see Refs.~[\\ref{Martel:1997vi:X}~--~\\ref{Bousso:2009ks:X}].} We fix other physical parameters and ask what values of the cosmological constant $\\Lambda$ we are likely to observe.  In Section~\\ref{sec:measure} we begin by reviewing the proposal of Ref.~\\cite{Nomura:2011dt}, and we then explain how cosmological parameters can be calculated in Section~\\ref{sec:prediction}.  While the framework itself is well-defined, any practical calculation is necessarily approximate, since we need to model ``experimenters'' who actually make observations.  In our context, we need to consider galaxy formation and life evolution therein, which we will do in Section~\\ref{sec:calc}.  We present the result of our calculation in Section~\\ref{sec:cc}.  We find that, in contrast to the case with some earlier measures~\\cite{Salem:2009eh}, the measure of Ref.~\\cite{Nomura:2011dt} does not lead to unwanted preference for a negative cosmological constant---in fact, a positive value is preferred. We find that a simple anthropic condition based on metallicity of stars is sufficient to make the calculated distribution consistent with the observed value at an order of magnitude level.  We conclude in Section~\\ref{sec:concl}.  Appendix~\\ref{app:calc} lists formulae for galaxy formation used in our analysis.  Appendix~\\ref{app:metal} discusses the anthropic condition coming from metallicity of stars. ",
        "conclusions": "\\label{sec:concl}  In this paper, we have studied the probability distribution of the cosmological constant (or the vacuum energy $\\rho_\\Lambda$) in the multiverse, using the quantum measure proposed in Ref.~\\cite{Nomura:2011dt}. We have found that this measure does not lead to a strong preference for negative $\\rho_\\Lambda$, as opposed to earlier measures proposed based on geometric cutoffs, because it does not experience a large volume effect associated with the global geometry of anti-de~Sitter space. Moreover, we have found that a positive value of $\\rho_\\Lambda$ is preferred, consistent with observation.  We have found that a simple, intuitive condition based on metallicity is enough to reproduce the observed value of $\\rho_\\Lambda$ at an order of magnitude level.  This is comfortable because effects from other possible anthropic conditions, such as the ones from encounters, are much more sensitive to the details of the conditions. \\begin{figure}[t] \\center{\\includegraphics[scale=0.7]{figure/Shading.eps}} \\caption{The normalized probability distribution $P(\\rho_\\Lambda)$ with a metallicity condition:\\ Eq.~(\\ref{eq:n-1}) with $m=2$. The light and dark shaded regions indicate those between $1$ and $2\\sigma$, and outside $2\\sigma$, respectively.  The observed value $\\rho_\\Lambda/\\rho_{\\Lambda,{\\rm obs}} = 1$ (denoted by a vertical line) is consistent with the distribution at the $1\\sigma$ level.} \\label{fig:final} \\end{figure} In Fig.~\\ref{fig:final}, we present the normalized distribution $P(\\rho_\\Lambda)$ with the $m=2$ metallicity constraint, where the $1$ and $2\\sigma$ regions are indicated.  We find that the observed value is consistent with the calculated distribution at the $1\\sigma$ level.  It would be interesting to refine our analysis including more detailed anthropic effects, such as those of star formation.  Another possible extension of the analysis is to vary other cosmological parameters, such as the primordial density contrast $Q$ and spatial curvature $\\Omega_k$ (at a specified time), in addition to $\\rho_\\Lambda$.  We plan to study these issues in the future."
    },
    "1107/1107.0448_arXiv.txt": {
        "abstract": "{ We present MERLIN observations of OH, water and methanol masers towards the young high mass stellar object IRAS~$20126+4104$. Emission from the 1665-MHz OH, 22-GHz H$_2$O and 6.7-GHz CH$_3$OH masers is detected and all originates very close to the central source. The OH and methanol masers appear to trace part of the circumstellar disk around the central source.  The positions and velocities of the OH and \\methanol\\ masers are consistent with Keplerian rotation around a central mass of ${\\sim}$ 5~\\Msun. The water masers are offset from the OH and \\methanol\\ masers and have significantly changed since they were last observed, but still appear to be associated to the outflow from the source.  All the OH masers components are circularly polarised, in some cases reaching 100 percent while some OH components also have linear polarisation. We identify one Zeeman pair of OH masers and the splitting of this pair indicates a magnetic field of strength $\\sim11$ mG within $\\sim0.5''$ (850 AU) of the central source. The OH and \\methanol\\ maser emission suggest that the disk material is dense, $n>10^6$ cm$^{-3}$, and warm, $T>125$K and the high abundance of \\methanol\\ required by the maser emission is consistent with the evaporation of the mantles on dust grains in the disk as a result of heating or shocking of the disk material. ",
        "introduction": "The extreme brightness and compact size of masers make them valuable probes of regions of high mass star formation.  However it is not clear how, or indeed if, the different types of maser, \\htwoo\\ , OH and \\methanol, are related to each other towards typical high mass protostars.  Studies of \\htwoo\\ and \\methanol\\ masers towards a number of sources have failed to find a coherent view of their relationship (e.g. Beuther et al 2002), however statistical studies of OH and \\methanol\\ masers suggest that their associations can tell us about the evolutionary sequence of star formation (e.g.  Caswell 1997; Szymczak, \\& Gerard 2004).  While the presence of particular types of maser towards a source may trace the evolutionary stage of the source, many sources show emission from multiple types of masers. Since the different masers can require different excitation conditions, these multiple species can be used as high resolution probes of different components of the circumstellar environment. Theoretical models suggest that the spatial coincidence of different maser transitions or maser species can be used to infer the properties of the emitting material. For example, Cragg et al. (2002) found that gas phase molecular abundance is the key determinant of observable maser activity for both OH and \\methanol\\ molecules. Detailed modelling of high resolution observations towards particular sources can also provide value insights on the structure of the circumstellar material. For example, models have shown that the complexity of observed magnetic field structure in W75N can be explained by the maser emission originating from different depths within the protostellar disk (Gray et al. 2003).  With these possibilities in mind, we present here the results of a study of the \\htwoo\\ , OH and \\methanol\\ masers associated with the well studied high mass young source IRAS~$20126+4104$. This source is located in the Cygnus X region at an estimated distance of 1.7 kpc (Wilking et. al. 1989).  It has a luminosity of $10^4$ \\Lsun\\ and is perhaps the best studied example of massive protostar associated with a Keplerian disk and a jet/outflow system. The molecular outflow has been mapped by Cesaroni et al. (1997; C97) and observations at centimetre wavelengths have shown that the outflow is fed by a jet (Hofner et al. 1999). On the other hand, CH$_3$CN (5-4) observations (C97) have revealed a molecular disk almost perpendicular to the jet axis and rotating around the embedded young stellar object (YSO) at the origin of the outflow/jet. Subsequent observations in the CH$_3$CN (12-11) (Cesaroni et al. 1999; C99) and NH3 (1,1) lines (Zhang et al. 1998) have found evidence for Keplerian rotation, implying a central disk plus stellar mass of 24 \\Msun.  The source is associated with OH, \\htwoo\\ and \\methanol\\ masers. Two features of OH masers were detected in the 1665-MHz line by Cohen et al. (1988) and more recent unpublished VLA data (Cohen, priv. comm.). Observations of the water masers using the VLA with angular resolution of $0.1''$ identified three emission regions (Tofani et al. 1995). Moscadelli et al. (2000; hereafter MCR) resolved two of these into 26 unresolved spots using the VLBA. The velocity and spatial structure of these spots were well fitted by a model with the spots arising at the interface between a jet and the surrounding molecular gas. Maser emission from the 6.7-GHz line of \\methanol\\ has recently been observed by Minier et al. (2001). Using the EVN, they found two maser clumps separated by 100 AU and $0.8''$ to the north of the central source, a region relatively remote from the central source and with no other known indication of activity related to star formation.  To determine how these three types of masers are related to the circumstellar disk and/or the jet and investigate the connection between them, we have observed IRAS~$20126+4104$ at high angular resolution using MERLIN.  The details of the observations and reduction are given in Sec. 2 and the results presented in Sec. 3. In Sec. 4 we discuss the interpretation while conclusions are drawn in Sec. 5.  \\begin{table*} \\begin{tabular}{|c|c|c|c|} \\hline  {Type of maser}  & {OH} & {H$_2$O} &{CH$_3$OH} \\\\ \\hline Date of observation & 20  & 28 \\& 29  & 16 \\& 17  \\\\ & Jan 2002 & Mar 2002 & May 2002 \\\\ Antenna Used & De CA DA  & CA MK2 DA  & CA MK2 \\\\ & KN MK2 TA & KN TA       &       \\\\ Field centre (2000) & $\\alpha  20^{h} 14^{m} 26.04^{s}$ & $\\alpha  20^{h} 14^{m} 26.04^{s}$ & $\\alpha  20^{h} 14^{m} 26.04^{s}$ \\\\ &$\\delta  41^{\\circ} 13' 32.5''$ & $\\delta  41^{\\circ} 13' 32.5''$ & $\\delta  41^{\\circ} 13' 32.5''$ \\\\ Rest frequency (MHz) & 1665.402 1667.359 & 22235.0798  & 6668.518  \\\\ No. of frequency channels & 1024 & 256 & 512 \\\\ Total band width (MHz) & 1 & 4 & 2 \\\\ % bandpass calibrator & 3C84 & 4C39.25 & 3C84 \\\\ Polarisation angle calibrator & 3C286 & No & No \\\\ Phase calibrator & 2005+403 & 2005+403 & 2005+403 \\\\ \\hline \\end{tabular} \\caption{Observing and calibration parameters for the MERLIN spectral-line observations of IRAS~$20126+4104$} \\end{table*} ",
        "conclusions": "We have used MERLIN to study the immediate vicinity of IRAS~$20126+4104$ at high angular resolution and have shown that the 1665 MHz OH, \\htwoo\\ and \\methanol\\ masers towards this source all originate within $\\sim0.5''$ (850 AU) of the central source. The OH masers have an elongated distribution, tracing part of a disk of material around the source which is orthogonal to the axis of the jet from the soure.  We could identify one Zeeman pair of OH masers which indicates a magnetic field of strength $\\sim$11 mG in this disk. The velocity structure of the OH masers is consistent with Keplerian motion around a central source of $\\stackrel{<}{\\sim}$10~\\Msun. The methanol masers are intermingled with the north-western part of the OH maser distribution and are at velocities intermediate between the north-western OH masers and those to the south-east.  Our observations confirm the close association of OH and methanol masers. We suggest that the high methanol (and OH) column densities necessary for the maser emission may result from the release of mantles from the dust grains in the surface layers of the circumstellar disk as the disk material has been heated by the central young star or as the stellar wind has shocked the disk material.  The \\htwoo\\ masers have significantly varied since they were last observed at high angular resolution. We detect only one of the three clusters previously seen. Although the \\htwoo\\ masers detected are close to the location central source, as was seen in the previous observation, the maser spots have a considerably different spatial and kinematic structure to those previously measured. In particular the detailed model proposed by MCR for the \\htwoo\\ masers arising at the survey an outflow cavity fails to acount for the velocities and locations of the current spots.  These observations show that, at least for this source, the three common types of maser associated with young high mass stars probe different components of the circumstellar environment allowing a coherent view of the circumstellar regions to be constructed. The OH masers provide a measurement of the magnetic field in the circumstellar disk within $\\sim500$AU of a young high mass star.  \\vspace{1cm} {\\bf ACKNOWLEDGMENTS}  We thank Peter Hofner for providing us with the 3.6cm map, and Riccardo Cesaroni, Malcolm Gray and Anita Richards for useful communication. MERLIN is a national facility operated by the University of Manchester at Jodrell Bank Observatory on behalf of PPARC."
    },
    "1107/1107.5601_arXiv.txt": {
        "abstract": "We measured temperatures, gravities, and masses for a large sample of blue horizontal branch stars in $\\omega$\\,Centauri, comparing the results with theoretical expectations for canonical and He-enriched stars, and with previous measurements in three other clusters. The measured gravities of $\\omega$\\,Cen stars are systematically lower than canonical models, in agreement with expectations for He-enriched stars, and contrary to that observed in the comparison clusters. However, the derived masses are unrealistically too low as well. This cannot be explained by low gravities alone, nor by any of the other parameters entering in the calculation. We find that the same stars are not brighter than their analogs in the other clusters, contrary to the expectations of the He-enrichment scenario. The interpretation of the results is thus not straightforward, but they reveal an intrinsic, physical difference between HB stars in $\\omega$\\,Cen and in the three comparison clusters. ",
        "introduction": "\\label{s_intro}  The nature of the second parameter, aside from metallicity, which determines the morphology of the horizontal branch (HB) in globular clusters (GCs), is one of the most longstanding problems of modern astrophysics. In fact, a lower metallicity favors the formation of hotter and bluer HB stars, but clusters with the same metallicity can show very different HB morphology \\citep{Sandage67}, and an HB extended far toward the blue is observed even in some metal-rich GCs \\citep[e.g.,][]{Rich97}. The helium abundance has been early proposed, among others, as this second parameter (\\citealt{Sweigart97,DAntona02}; see \\citealt{Catelan09b}, for a review) because, during the He-burning phase, helium-rich stars are expected to be hotter than objects of canonical composition. This model has recently drawn much attention, triggered by the discovery of multiple stellar populations in GCs \\citep{Piotto05,Piotto07}. In fact, \\citet{Piotto05} showed that a different metallicity is not the cause of the main-sequence split observed in $\\omega$\\,Centauri \\citep{Bedin04}, and the only explanation is that the bluer sequence is greatly enriched in helium, about 50\\% more He-rich (Y=0.38) than in normal metal-poor GC stars \\citep{Norris04,Piotto05}. In this scenario, the blue HB stars observed in many GCs could be the progeny of the He-enriched second stellar generation. Unfortunately, diffusion processes completely alter the surface chemical abundances of hot HB stars \\citep[e.g.,][]{Behr99}, preventing a direct demonstration of the connection between multiple populations and HB morphology. Nevertheless, an increased helium content can be indirectly deduced from other observable quantities, because He-enriched HB stars are predicted to be brighter \\citep{Sweigart87} and to occupy different loci in the temperature--gravity plane \\citep{Moehler03}.  In this Letter, we present the results of our investigation aimed to search for an indirect indication of helium enrichment among blue HB stars in $\\omega$\\,Centauri, to test the He-enrichment scenario and its predicted effects on the HB morphology. This cluster is the ideal target for our purpose, because it hosts a very complex stellar population, comprising three known MSs and six sub-giant branches \\citep{Bellini10}. ",
        "conclusions": "\\label{s_conclusions}  Blue HB stars in $\\omega$\\,Cen show lower gravities with respect to both canonical models and analogous stars in other GCs, but their stellar masses are underestimated, and their visual absolute magnitudes are very similar to those of the comparison clusters. Neither the low gravities nor the other parameters involved in the calculation can explain the too low masses. We can firmly conclude that these results reveal an intrinsic difference between the blue HB stars in $\\omega$\\,Cen and their analogs in other GCs, but its interpretation is not straightforward. The lower gravities follow the expectations for He-rich stars, but the magnitudes and masses do not."
    },
    "1107/1107.0212_arXiv.txt": {
        "abstract": "We performed a statistical analysis of half-power pulse--widths of the core components in average pulsar profiles. We confirmed an existence of the lower bound of the distribution of half-power pulse--width versus the pulsar period $W_{50}\\sim2^{\\circ}.45 \\; P^{-0.5}$ found by Rankin (1990). Using our much larger database we found $W_{50}=(2^{\\circ}.51\\pm0^{\\circ}.08) \\; P^{-0.50\\pm0.02}$ for 21 pulsars with double--pole interpulses for which measurement of the core component width was possible. On the other hand, all single--pole interpulse cases lie in the swarm of pulsars above the boundary line. Using the Monte Carlo simulations based on exact geometrical calculations we found that the Rankin's method of estimation of the inclination angle $\\alpha\\approx asin(2^{\\circ}.45\\; P^{-0.5}/W_{50})$ in pulsars with core components is quite good an approximation, except for very small angles $\\alpha$ in almost aligned rotators. ",
        "introduction": "It is generally accepted that the elements of mean pulsar profiles can be divided into core and conal components (\\citet{r83}, Rankin (1990, hereafter \\citet{r90}), \\citet{r93}). The core component usually occupies central parts of the profile and is often flanked by one or two pairs of conal components. The importance of the conal profile components was recently studied statistically by Maciesiak, Gil \\& Ribeiro (2011; hereafter \\citet{mgr11}). These authors used recently large data sets of pulsar periods, profile widths and inclination angles resulted from large pulsar surveys conducted in recent years. They also compiled the largest ever database of pulsars with interpulses (IP), divided into the double--pole (DP--IP) and single--pole (SP--IP) cases (see Table \\ref{tab.1} in this paper). The IP pulsars constitute about 3 percent (2 and 1 percent for DP--IP and SP--IP cases, respectively) of the pulsar population. The DP--IP pulsars (almost orthogonal rotators) tend to be younger or middle aged, while SP--IP pulsars tend to be older (see Figure 6 in \\citet{mgr11}). This suggest a secular alignment of the magnetic axis towards the spin axis over a timescale of about $10^7$ years.  The interpulse emission is important also for the statistics of the core profile components in pulsars. Rankin (1990; hereafter R90) measured the half-power (50\\% of the maximum intensity) widths of core components in a number of pulsars (about 100), including those showing the interpulse emission. She found that the distribution of pulse--widths has a clear lower bound $2^{\\circ}.45 \\; P^{-0.5}$ populated by DP--IP cases, in which the inclination angle $\\alpha$ must be close to $90^{\\circ}$. The rest of pulsars including SP--IP cases lied in a swarm of data points above the lower bound. \\citet{r90} described this distribution by a simple mathematical expression \\begin{equation} W_{50} = 2^{\\circ}.45 \\; P^{-0.5} / sin \\alpha, \\label{eq.1} \\end{equation} indicating that the core component widths depend only upon the pulsar period $P$ and the inclination angle $\\alpha$ between the rotation and magnetic axes of the neutron star and they are almost insensitive to the impact angle $\\beta$ of the closest approach of the observer's line--of--sight to the magnetic axis (for description of pulsar geometry see Paper I or Figure \\ref{figure.b1} in the on-line Appendix \\ref{appendix.b}). The above equation represents a useful method of estimation of the inclination angle $\\alpha$ in pulsars with core component in the mean profile. In this paper we attempt to revisit and verify these important and intriguing results, using a more abundant database that we compiled in \\citet{mgr11}). ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics{fig8.eps} \\caption{Plot of pulse--width $W_{50}$ versus period $P$ for 1495 normal pulsars from the ATNF database. \\label{figure.8}} \\end{center} \\end{figure} The most important result of our paper is a clear confirmation of the lower limit in the distribution of half-power pulse--width $W_{50}$ in core components of pulsar profiles when presented versus the pulsar period $P$ using much bigger database than that used originally by \\citet{r90}. As illustrated in Figures \\ref{figure.1}, \\ref{figure.2}, \\ref{figure.3} and \\ref{figure.8} this limit is represented by the line $W_{50}=2^{\\circ}.5\\: P^{-0.5}$ at 1 GHz and $W_{50}=2^{\\circ}.37\\: P^{-0.5}$ at 1.4 GHz. It seems that more appropriate description of the lower boundary limit is the lower boundary belt (LBB) $W_{50}=(2^{\\circ}.37^{+1^{\\circ}.43}_{-0^{\\circ}.87})\\: P^{-0.5}$, which accounts for scatter and  errors in $W_{50}$ measurements. Strictly speaking, the widths of the LBB was determined by the scatter and errors of DP-IP cases marked as red symbols in Figure 3. However, the normal data points seem to populate the belt marked in Figure 8 in similar manner as the DP-IP cases do in Figure 3. The concept of LBB is based not only on core widths of DP--IP pulsars (like in \\citet{r90}) but here it is based on all $W_{50}$ measurements of normal pulsars available in the ATNF database, without distinctions into core and conal components, of course. This is dramatically illustrated in Figure \\ref{figure.8} containing 1495 data points. This figure is a reproduction of Figure \\ref{figure.3} with all special points described in the legend omitted. Apparently, the existence of the lower boundary limit in the pulse--width measurements discovered by Rankin (1990) does not depend neither on the method nor on the errors of measurements of the core widths in DP-IP pulsars. The unweighted formal fit to the points within LBB (between dashed lines) is $W_{50}=(2^{\\circ}.86\\pm0^{\\circ}.07)\\: P^{-0.51\\pm0.02}$.  The existence of the lower bound in the distribution of $W_{50}$ versus pulsar period $P$ is very intriguing and requires solid and  convincing explanation. The exponent $-0.5$ is evidently related to the geometry of dipolar field lines in the radio emission region, while the value of the numerical factor about $2^{\\circ}.5$ is a little more mysterious. \\citet{r90} argued that the core component widths are intimately related to the polar cap geometry at the stellar surface, whose bundle of the last open dipolar field lines have an angular extent (opening angle $2\\rho$) very close to $2^{\\circ}.45 \\: P^{-1/2}$. Although this natural and very appealing explanation is commonly accepted, it is not free from theoretical and interpretational problems. First of all, the coherent radio emission cannot originate at or very near the polar cap, because there are no plasma instabilities available in this region. Recently, a modern version of this idea is discussed, with arguments that the core emission originates below the conal one but the former is emitted well above the polar cap surface. For example, \\citet{mitra07} argued that in PSR B$0329+54$ the emission of its core component originates over a range of altitudes (up to about 100 km) above the star surface (see also \\citet{gil91}). This broadens the required opening angle by a factor of about square root of the normalised emission altitude, that is about 3. If this is a case then one has to conclude that the core radio emission is related to an innermost subset of the open dipolar magnetic field lines. The problem therefore is to identify structures within the overall pulsar beam with an angular extent corresponding to the observed lower boundary limit. We plan to devote to this problem the forthcoming Paper III of this series.  The interpretational problems mentioned above do not affect the approximate method of estimating the inclination angle $\\alpha$ in any pulsar with a core component proposed by \\citet{r90}. We have shown by means of Monte Carlo simulations based on the strict geometrical calculations that this method offers quite good an approximation with the accuracy of 5, 10, 20 and 50 percent for $\\alpha\\sim90^{\\circ}$, $\\alpha\\sim50^{\\circ}$, $\\alpha\\sim25^{\\circ}$ and $\\alpha\\sim10^{\\circ}$, respectively (Figure \\ref{figure.7}). This method cannot be used in almost aligned cases with very small values of $\\alpha$."
    },
    "1107/1107.3161_arXiv.txt": {
        "abstract": "Core-accretion planet formation begins in protoplanetary disks with the growth of small, ISM dust grains into larger particles.  The progress of grain growth, which can be quantified using 10$\\micron$ silicate spectroscopy, has broad implications for the final products of planet formation.  Previous studies have attempted to correlate stellar and disk properties with the 10$\\micron$ silicate feature in an effort to determine which stars are efficient at grain growth.  Thus far there does not appear to be a dominant correlated parameter.  In this paper, we use spatially resolved adaptive optics spectroscopy of 9 T Tauri binaries as tight as 0.25\" to determine if basic properties shared between binary stars, such as age, composition, and formation history, have an effect on dust grain evolution.  We find with 90-95\\% confidence that the silicate feature equivalent widths of binaries are more similar than those of randomly paired single stars, implying that shared properties do play an important role in dust grain evolution.  At lower statistical significance, we find with 82\\% confidence that the secondary has a more prominent silicate emission feature (i.e., smaller grains) than the primary.  If confirmed by larger surveys, this would imply that spectral type and/or binarity are important factors in dust grain evolution. ",
        "introduction": "}  The core-accretion model for planet formation \\citep[e.g.,][]{2007prpl.conf..591L} is a multi-step process beginning with the agglomeration of small interstellar medium (ISM) dust grains into larger particles.  Eventually, these particles reach the size of planetesimals and gravitationally attract each other and their surrounding gas.  Planet formation must occur quickly, as several physical processes are able to disperse the gas and dust of the protoplanetary disk.  Consequently, the timescale over which each aspect of core-accretion occurs is critical to the final outcome of the planet formation process.  The first step of core-accretion, where small ISM dust grains agglomerate into larger particles, is of particular importance as it sets the conditions for all of the other core-accretion steps.  Dust grain properties can be studied at a variety of wavelengths and locations in the disk \\citep{2007prpl.conf..767N}.  Visible and near-infrared scattered light imaging can be used to study the grain sizes of submicron and micron surface dust at large radial distances.  Mid-infrared spectroscopy of the 10$\\micron$ and 20$\\micron$ silicate emission features can be used to determine the size and composition of submicron and micron sized surface dust at small radial distances.  Far-infrared and millimeter continuum observations probe deep into the disk mid-plane over the full radius of the disk and can be used to determine the size distribution and total mass of mm-cm dust grains.  Although these observations probe different grain sizes at different locations in the disk, it appears that they are somewhat correlated.  In particular, the equivalent width of the 10$\\micron$ silicate feature is correlated with the slope of the submillimeter spectral energy distribution (SED), implying that dust grain growth proceeds concurrently through its various phases \\citep{2010A&A...515A..77L}.  Additionally, the shape of the 10$\\micron$ silicate feature is correlated with its amplitude, likely due to the simultaneous development of large amorphous grains and crystalline grains \\citep{2003A&A...400L..21V,2003A&A...412L..43P,2005Sci...310..834A}.  Our current understanding of dust grain growth is that it begins during the Class II phase of star/planet formation, once the dust is collected into a circumstellar disk \\citep{1991ApJ...381..250B,1994A&A...291..943O,1994A&A...288..929K}.  At this point, dust grains are expected to coagulate and settle rapidly \\citep[$\\sim10^{3}-10^{6}$ years depending on model assumptions;][]{2005AA...434..971D,2008AA...480..859B}.  However, Class II objects over a wide span of ages ($\\sim$0.5-10 Gyr) have a diverse set of dust grain properties, implying that dust evolution is not simply a function of age \\citep{2001A&A...365..476M,2003A&A...412L..43P,2005A&A...437..189V,2005Sci...310..834A,2006ApJ...639..275K,2009ApJ...703.1964F,2011arXiv1104.3574O}.  Even within individual clusters, which are approximately coeval, there is no apparent correlation between mid-infrared dust properties and stellar mass, luminosity, stellar accretion or disk mass \\citep{2007ApJ...659.1637S,Watson2009}.  Over a large range of masses, taken between stars in different clusters with different ages, there appears to be a weak correlation between mass and silicate feature strength \\citep{2006ApJ...639..275K,2009ApJ...696..143P}, which might be the result of another spectral-type dependent factor, such as X-rays or the location of the silicate emission zone \\citep{2009A&A...508..247G,2007ApJ...659..680K}.  However, the general diversity of dust properties between similar stars of similar ages suggests that additional properties are needed to explain the evolution of the dust.  One way to isolate which stellar properties are important for dust grain evolution is to observe the dust grain properties of coeval young binaries.  Binaries share certain properties, such as age, composition and formation history, which are difficult to ascertain in individual stars.  If these properties are important for dust grain evolution, then we expect the dust to be similar in binary pairs.  However, if these properties are unimportant, then the binary stars will have a random distribution of dust characteristics, similar to those observed in single stars.  Additionally, if the shared properties are important, they might be hiding correlations between other stellar parameters and dust grain characteristics, which can be uncovered by taking advantage of the coevality of binaries.  In this paper \\citep[and including the results of ][]{Skemer2010}, we present spatially resolved 10$\\micron$ silicate spectroscopy of 8 Taurus-Aurigae\\footnote{Taurus-Auriga is a low-mass star forming region with an age of $\\sim$1-2 Myr and a distance of $\\sim$140 pc \\citep{1995ApJS..101..117K,1994AJ....108.1872K}.} binaries, which triples the sample size of spatially/spectrally resolved binaries in a single cluster at mid-infrared wavelengths.  Small diameter space-based telescopes are unable to perform high spatial resolution ($\\lesssim 2-3$\") observations in the mid-infrared ($>$8$\\micron$) due to their large diffraction limits (additionally, the Infrared Spectrograph, IRS, on \\textit{Spitzer} is no longer operational, so mid-infrared spectroscopy is not currently accessible from space).  Ground-based observatories have much larger diameters but suffer from high sky background, variable transmission and seeing, which explains why constructing even a small sample has been difficult.  We describe our observations and reductions in Section 2 and our corrections for extinction and equivalent width measurements in Section 3.  In Section 4, we perform statistical tests in an attempt to determine which properties affect dust grain-growth.  In particular, we ask the questions (1) \\textit{Is grain growth correlated between binary pairs?} and (2) \\textit{Does removing this effect reveal correlations between grain growth and other properties?}  We discuss the implications of our results and our conclusions in Section 5. ",
        "conclusions": "Our mid-infrared adaptive optics survey has produced 8 new spatially resolved silicate spectra of young Taurus-Aurigae binaries \\citep[including UY Aur, which was published in][]{Skemer2010}.  This effectively triples the number of spatially and spectrally resolved binaries in the Taurus-Auriga cluster, and provides the first opportunity to draw statistical conclusions about the spatially resolved silicate features of binaries in any young cluster.  Our fundamental questions are (1) \\textit{Is grain growth correlated between binary pairs?} and (2) \\textit{Does removing this effect reveal correlations between grain growth and other properties?}  We find with 90\\% confidence (when including GG Tau) or 95\\% confidence (when excluding GG Tau) that the silicate features of binaries are more similar than the silicate features of randomly paired single stars (our ambivalence in including GG Tau stems from its large circumbinary disk, which might be replenishing the dust in the inner disks via streamers).  The similarity of the silicate features in our binary pairs implies that one or more shared binary properties (such as age, composition or formation history) plays an important role in dust grain evolution.  One possible explanation for our findings is that the environment in which a star forms affects its potential to grow dust grains, and eventually planets.  It is a fundamental prediction of core-accretion that giant planets form more efficiently in metal-rich systems, based on the increased solid material available for building planetary cores \\citep{2004ApJ...604..388I,2004ApJ...616..567I}.  However, the core-accretion scenario is predicated by small dust grains coagulating quickly enough that aerodynamic drag does not remove too many meter-sized objects from the system \\citep[a problem, which is commonly known as the ``meter-sized barrier\";][]{1977MNRAS.180...57W}.  \\citet{2008AA...480..859B} have shown that the removal of meter-sized objects due to radial drift can be overcome by increasing the disk's initial dust-to-gas ratio from 1\\% to 2\\%.  The result is a natural consequence of increasing the rate of grain growth by decreasing the time between collisions.  Since binary stars form from the same fragmented core material, it is possible that grain growth in binary stars will proceed at similar rates in each component due to their shared initial gas-to-dust ratio.  Binarity itself might also affect the silicate features of young stars.  Circumstellar disks can be dynamically perturbed by a companion star, which introduces spiral structure and truncates the radial extent of the disk \\citep{1994ApJ...421..651A}.  In this configuration, large particles collide and fragment more frequently, which speeds up dust grain evolution \\citep{1995Icar..114..237D} while decreasing the maximum particle size in the disk \\citep{2011A&A...527A..10Z}.  Dust streamers from a wider circumbinary disk can also replenish solid material in one or both of the circumstellar disks \\citep{1996ApJ...467L..77A}.  The dust streamers are likely to contain un-evolved, ISM-like dust grains that dominate the appearance of the silicate feature \\citep{Skemer2010}.  There are two young Taurus-Aurigae binaries with directly imaged circumbinary disks: GG Tau \\citep{1994A&A...286..149D} and UY Aur \\citep{1996A&A...309..493D}.  In each case, we find that one of the components has an unusually prominent silicate feature, indicating the presence of small grains.  \\citet{2008ApJ...673..477P} were not able to detect a difference between the silicate features of (unresolved) medium-separation binaries and single stars, which suggests that any effect must be subtle or uncommon.  However, spatially resolved studies of the silicate features in young binaries might reveal different dust-grain growth mechanisms.  In Section \\ref{Binaries others?}, we found that seven of our nine binaries (i.e. 82\\% confidence of a trend using the binomial distribution) have larger silicate equivalent widths (smaller grains) in the secondary than the primary.  This pattern was also observed in two binaries by \\citet{2003A&A...412L..43P}.  Assuming this trend is confirmed by larger samples of spatially resolved silicate spectroscopy, the cause could be based on the dynamics of binaries or simply a spectral-type/mass/luminosity correlation with the speed of dust grain growth.  The former explanation could be tested by seeing if the correlation is connected with binary separation.  Our study has been limited by its small sample size, due to the complexity and limitations of ground-based 10 $\\micron$ spectroscopy and adaptive optics, which were necessary to build up the current sample.  Larger ground-based telescopes, such as the LBT and ELTs along with large space-based telescopes, such as JWST will be able to dramatically increase the sample of spatially resolved silicate spectra of young binaries."
    },
    "1107/1107.4094_arXiv.txt": {
        "abstract": "We investigate the non-linear evolution of the matter power spectrum by using a large set of high-resolution N-body/hydrodynamic simulations.  The linear matter power in the initial conditions is consistently modified to mimic the presence of warm dark matter particles which induce a small scale cut-off in the power as compared to standard cold dark matter scenarios. The impact of such thermal relics is examined at small scales $k>1\\,\\ihmpc$, at redshifts of $z<5$, which are particularly important for the next generation of \\lya forest, weak lensing and galaxy clustering surveys.  We measure the mass and redshift dependence of the warm dark matter non-linear matter power and provide a fitting formula which is accurate at the $\\sim$ 2\\% level below $z=3$ and for particle masses of \\mwdm $\\ge$ 0.5 keV. The role of baryonic physics on the warm dark matter induced suppression is also quantified. In particular, we examine the effects of cooling, star formation and feedback from strong galactic winds. Finally, we find that a modified version of the halo model describes the shape of the warm dark matter suppressed power spectra better than {\\small{HALOFIT}. In the case of weak lensing however, the latter works better than the former, since it is more accurate on the relevant, mid-range scales, albeit very inaccurate on the smallest scales ($k>10\\,\\ihmpc$) of the matter power spectrum.} ",
        "introduction": "The increasing amount of observational data available and the numerical tools developed for their interpretation have given rise the so-called era of precision cosmology. At the present time, the cosmological concordance model based on a mixture of cold dark matter and a cosmological constant must thereby be tested in new regimes (both in space and in time) and using as many observations and techniques as possible in order to either confirm or disprove it.  Among the many different observables the non-linear matter power spectrum is a crucial ingredient since it allows us to describe the clustering properties of matter at small scales and low-redshift, where linear theory is not reliable. However, accurate modelling of non-linear physical processes is needed, in order to use this observable to gain quantitative results on the nature of dark matter.  Warm Dark Matter (WDM) is an intriguing possibility for a dark matter candidate. Its velocity dispersions are intermediate between those of cold dark matter and hot dark matter (e.g. light neutrinos). In this scenario, at scales smaller than WDM free-streaming scales, cosmological perturbations are erased and gravitational clustering is significantly suppressed. If such particles are initially in thermal equilibrium, they have a smaller temperature and affect smaller scales than those affected by neutrinos. In addition, WDM produces a distinctive suppression feature at such scales as compared to that induced by neutrinos. For example, thermal relics of masses at around 1 keV which constitute all of the dark matter have a free-streaming scale that is comparable to that of galaxies, well into the non-linear regime.  Among the different WDM candidates a special role is played by the sterile neutrino with mass at the keV scale (\\cite{boya09a}). WDM was originally proposed to solve some putative problems that are present in cold dark matter scenarios at small scales (see \\cite{colin00,bode}), however it is at present controversial whether these tensions with cold dark matter predictions can be solved by modifying the nature of dark matter particles or by some other baryonic process (e.g. \\cite{trujillo10}).  In the present paper we wish to quantify the impact of a WDM relic on the non-linear power spectrum by using a set of N-body/hydrodynamic simulations of cosmological volumes at high resolution. Investigating WDM scenarios in a cosmological setting has been done by means of N-body codes in order to carefully quantify the impact of such a candidate in terms of halo mass function, structure formation, halo density properties (\\cite{bode,colin08,colombi09}) and particular care needs to be placed on correctly addressing numerical/convergence issues (\\cite{ww07}). In general, while the WDM induced suppression transfer function can be reliably estimated in the linear regime (e.g. \\cite{viel05,boya08,les11}), the non-linear suppression has not been investigated. However, a recent attempt to obtain the non-linear matter power at small scales by modifying the halo model is described in \\cite{smith11}.  The analysis of matter power spectra at small scales has been performed in recent year by different groups by focussing mostly on baryon physics such as feedback and cooling (e.g.  \\cite{rudd08,guillet10,casarini11,vandaalen11}).  Recently, \\cite{vandaalen11} presented an extensive investigation of the effects of several different implementations of galactic feedback on the total matter power. Including feedback from Active Galactic Nuclei (AGN) was claimed to be most realistic since it matches optical and X-ray observations of groups of galaxies and solves the overcooling problem (see also related works by \\cite{puchwein08,fabjan010,mccarthy10,teyssier11}). This scenario has a relatively large impact in terms of total matter power at the scales affected by the presence of a WDM candidate. In a subsequent paper, \\cite{semboloni11} carefully analysed the impact of this feedback on weak lensing observables and concluded that it will not be possible to constrain WDM masses with future weak lensing surveys (as claimed in for example \\cite{markovic11}). These findings are very interesting and show that a future measurement of the matter power at such small scales should be considered with an accurate model of baryonic physics and not only of the dark matter component. In the present work, although we will be exploring some feedback mechanism that do not include the AGN feedback, the main focus is on the signature of the WDM suppression in terms of total matter power, presented as a ratio with a corresponding $\\Lambda$CDM standard model that includes the same astrophysical input and differs only in the initial total matter power.  Different constraints can be obtained by using several astrophysical probes.  For example by using \\lya observables such as the transmitted \\lya flux power, very competitive measurements in the form of lower limits ($m_{\\rm WDM}>4$ keV, 2$\\sigma$ C.L.) have been derived by using the SDSS flux power and other higher redshift and higher resolution data \\citep{viel05,viel06wdm,seljak06wdm}: these constraints become much weaker if the WDM is assumed to account to only a given fraction of the dark matter \\citep{boya09a} or if the initial linear suppression for a sterile neutrino is considered \\citep{boya09b} (in this latter case basically any $m_{\\rm WDM,sterile}>1$ keV is allowed). Alternatively, constraints on WDM models can be placed using: the evolution and size of small scale structure in the local volume high resolution simulations \\citep{tikhonov09}; simulated Milky Way haloes to probe properties of satellite galaxies \\citep{polisensky11,lovell11}; large scale structure data \\citep{abazajian06}; the formation of the first stars and galaxies in high resolution simulations \\citep{gao07}; weak lensing power spectra and cross-spectra \\citep{markovic11,semboloni11}; the dynamics of the satellites \\citep{knebe08}; the abundance of sub-structures \\citep{colin00}; the inner properties of dwarf galaxies \\citep{strigari06}; mass function in the local group as determined from radio observations in HI \\citep{zavala09}; the clustering properties of galaxies at small scales \\citep{coil08}; the properties of satellites as inferred from semi-analytical models of galaxy formation \\citep{maccio010}; phase-space density constraints from dwarf galaxies \\citep{devega10}.    We believe that most of the astrophysical probes used so far in order to constrain the small scale properties of dark matter could benefit from a comprehensive numerical modelling of the non-linear matter power.  The present work aims at providing such a quantity by using N-body/hydrodynamic simulations. The findings could also be useful for future surveys such as PanSTARRS, DES, LSST, ADEPT, EUCLID, JDEM or eROSITA, WFXT and SPT.   The layout of the paper is as follows.  In Section 2 we present our set of simulations and the code we use in order to investigate the non-linear suppression on the total matter power. Section 3 contains the main results of the present work and the description of the checks made in order to present a reliable estimate of the WDM non-linear suppression: we focus on numerical convergence, box-size, baryonic physics, particle velocities and the effect induced by cosmological parameters on the WDM power. As an application of the findings in Section 3 we present the weak lensing power and cross-spectra for a realistic future weak lensing survey in Section 4 and compare these results with those that could be obtained by using either linear-theory or halo models in Section 5. We conclude with a summary in Section 6. ",
        "conclusions": "By using a large set of N-body and hydrodynamic simulations we have explored the non-linear evolution of the total matter power.  The focus of the present work is on small scales and relatively low redshifts where non-linear effects are important and need to be properly modelled with simulations.  We checked for numerical convergence and box-sizes/resolution effects in the range $k=1-10$ \\ihmpc. We explored how different masses of a warm dark matter candidate affect the non-linear suppression as compared to a corresponding $\\Lambda$CDM model that shares the same parameters and astrophysical inputs. Our findings can be summarized as follows: \\begin{itemize} \\item[-] Cosmological volumes of linear size 25$h^{-1}$ comoving Mpc and with $512^3$ DM particles are sufficient to sample the WDM suppression for \\mwdm $\\ge 1$ keV at the percent level at $k<10$ \\ihmpc. \\item[-] The non-linear suppression induced by WDM is strongly redshift dependent. However, by $z=0$, up to $k=10$ \\ihmpc, there are virtually no differences (below 1\\%) between $\\Lambda$CDM and WDM models with \\mwdm $\\ge$ 1 keV. \\item[-] At higher redshifts differences are larger, being closer to the linear suppression. At $z\\sim 1$ there are differences of the order of a few percent between the non-linear WDM and $\\Lambda$CDM power spectra. \\item[-] Baryonic physics and in particular radiative processes in the gas component, the star formation criterion and galactic feedback in the form of winds are likely to affect the matter power at the 2-3 \\% level in the range $k=1-10$ \\ihmpc. However, a much stronger effect can be expected if AGN feedback is considered as in (\\cite{vandaalen11}). \\item[-] We investigate how a change in $\\Omega_{\\rm m}$, $H_0$ and $\\sigma_8$ impacts the non-linear power and WDM suppression in particular, when values different from our reference choice are used. Small difference are found (at the $\\pm$ 2\\% level) at the scales considered here. Thus the WDM cutoff has a distinctive feature which is not degenerate with other cosmological parameters also at a non-linear level. \\item[-] We provide a useful fit to the non-linear WDM induced suppression in terms of a redshift-dependent transfer function; this fitting formula should agree to the actual measured power at the 2\\% level at $z<3$ and for masses above 0.5 keV. \\item[-] Reaching a higher accuracy (percent level) in terms of WDM non-linear power would require a very extensive analysis of astrophysical aspects related to the baryonic component such as considering different feedback effects. Among these the most promising seems to be AGN feedback which happen to solve the overcooling problem and has a strong impact on the total matter power (see \\cite{vandaalen11,semboloni11}). \\item[-] We find that future weak lensing surveys will most likely be powerful enough to rule out WDM masses smaller than 1 keV, which is consistent with previous results of \\cite{markovic11} and \\cite{smith11}. Ruling out models for masses larger than 1 keV would still be possible by using the cross-correlation signal between different redshift bins. However, measurement of the weak lensing power at these scales should also consider the effect of baryonic physics carefully and parameters could be biased as recently found in \\cite{semboloni11}, where it has been shown that AGN feedback produces a suppression which is larger than the one induced by WDM at the scales considered here. \\item[-] Non-linear corrections to the matter power spectrum in the WDM scenario obtained from {\\small{HALOFIT}} correspond better to the results of the WDM only simulation at scales $k<10$ \\ihmpc, if compared to the non-linear corrections of the halo model from Smith \\& Markovic (2011). Because these scales are most relevant for weak lensing power spectra, using {\\small{HALOFIT}} yields a better correspondence to the weak lensing power spectra calculated using our fitting function. However, on scales $k>10$ \\ihmpc, the halo model performs slightly better in that it better describes the shape of the suppression in the power spectrum, even if it does overestimate the effect. For this reason we believe that a further modification to the halo model may be needed, especially for weak lensing power spectra calculations. \\end{itemize}  As recently shown by \\cite{vandaalen11} and \\cite{semboloni11} including AGN feedback has strong consequences in terms of matter power and weak lensing, a comprehensive analysis that aims at measuring the mass of a warm dark matter candidate should thus hope to lift the degeneracies present (i.e. suppression in terms of matter power) by exploiting the different redshift and scale dependencies and by fully exploring the astrophysical parameter space and marginalize over the nuisance parameters.  We believe that future efforts aiming at measuring the coldness of cold dark matter should investigate the non-linear matter power in the range $z=0-5$ either using weak lensing observables or the small scale clustering of galaxies. These constraints can be particularly useful since they are complementary to those that can be obtained from high redshift \\lya forest data (e.g. BOSS/SDSS-III survey) or galactic and sub-galactic observables in the local universe."
    },
    "1107/1107.5484_arXiv.txt": {
        "abstract": "Galactic jets are powerful energy sources reheating the intra-cluster medium in galaxy clusters. Their crucial role in the cosmic puzzle, motivated by observations, has been established by a great number of numerical simulations missing the relativistic nature of these jets. We  present the first relativistic simulations of the very long term evolution of realistic galactic jets. Unexpectedly, our results show no buoyant bubbles,  but large  cocoon regions compatible with the observed X-ray cavities. The reheating is more efficient and faster than in previous scenarios, and it is produced by the shock wave driven by the jet, that survives for several hundreds of Myrs. Therefore, the X-ray cavities in clusters produced by powerful relativistic jets would remain confined by weak shocks for extremely long periods, whose detection could be an observational challenge. ",
        "introduction": "Galaxy clusters are formed by dark matter and gas. This last component appears in the form of galaxies and a diffuse hot gas filling the space amid them -- the intra cluster medium (ICM). The basic laws of physics predict that huge amounts of this gaseous component -- ICM -- should cool due to bremsstrahlung radiation and fall onto the central galaxies in the clusters. These flows of cold gas would eventually be intimately related with crucial processes in the galaxy formation, like for instance, the star formation history in galaxies. However, these flows -- the so called cooling flows -- have not been observed in most of the clusters, or when observed, they are not as important as expected. In order to reconcile the theoretical results with the observations, several physical mechanisms have been invoked, being the most widely accepted, the reheating of the ICM by the Active Galactic Nuclei (AGN) feedback \\citep{fab94,mn07}.  The state-of-the-art picture of the AGN feedback is underpinned on the idea that galactic jets can transport energy from the very center of the galaxy out to the cluster scales. These energy injections would inflate bubbles whose evolution would have two well differentiated phases: first, a shock dominated supersonic phase, followed by a second subsonic phase when the bubbles inflated by the jets would be  buoyant in the ICM \\citep{cat09}.  The buoyant bubbles are unstable, due to Rayleigh-Taylor and Kelvin-Helmoltz instabilities, when interacting with the surrounding ICM, leading to a mixture of the hot  gas locked in the bubble with the environment. Besides, an additional mixing is produced at the turbulent wake created by the buoyant bubbles rising up in the cluster potential well. All this mixing produces a net gain of internal energy of the ICM resulting in an efficient feedback mechanism able to stop or delay the cooling flows \\citep[see, e.g.,][]{ch01,qb01,br02,ch02,rb02,bi04,dv04,ry04,br06,ss08,dy10}. The idea that the radio lobes reach pressure equilibrium with their environment in a relatively short time supports the model of buoyant motion of the bubbles inflated by the jets \\citep{mn07}. However, the observations of shocks in several sources like  Hercules~A \\citep{nu05}, Hydra~A \\citep{sim09}, MS0735.6+7421 \\citep{mc05} or HCG~62 \\citep{git10} may require to reconsider the relevance of the subsonic buoyant phase. Specially, when observational evidences \\citep{kr07,cr07} and numerical simulations \\citep{pm07,bbrp11} have shown that even very modest energy injections -- low-power FRI jets -- with ages $\\simeq 10^7$ yrs still present relatively strong shocks.  The AGN feedback scenario has mainly been stablished by a great number of numerical simulations studying the long term evolution of jets \\citep[e.g.,][]{ch01,qb01,br02,rey02,om04,ob04,za05,br06,vr06,ct07,vr07,br07,bin07,bsh09,oj10}. In these works, the input of jets was modeled injecting mass, momentum and energy in a few computational cells with a huge size (for jet scales) and with low (i.e., non-relativistic) flow velocities and temperatures imposed by the Newtonian approach. These inherent constrains could have direct implications on the evolution of the simulated jets since, in order to match the typical momentum and energy fluxes, the jets are set up with unrealistic opening angles, radii, and flux masses when compared with observations. This could be the reason leading, in general, to the formation of weak shock waves (with Mach numbers $M_s\\leq 5$) and, as a consequence, to an early transition to the subsonic phase in almost all the simulations performed until now.  In this paper we present, for the first time, the results of a set of very long term axisymmetric FRII-like jet simulations evolved using a fully relativistic description of the fluid dynamics and thermodynamics, also including a relativistic equation of state. Our approach allows the simulations to reconcile the inferred momentum and energy fluxes of observed jets with reasonable values of the jet flow velocities, radii, opening angles and mass fluxes. The use of the relativistic equation of state accounts consistently for the relativistic character of electrons in the jet and, to a less extent, in the cocoon, and the non-relativistic behaviour of protons. Also our relativistic models correctly describe, by construction, the deceleration of the relativistic flow and the internal and kinetic energy fluxes accross the jet terminal shock, that govern the effective energy flux into the cocoon. ",
        "conclusions": "We have presented the first simulations on the impact of relativistic AGN jets in the heating of the intracluster medium. Our simulations cover the longest spatial and temporal scales considered up to now and besides this, due to their relativistic character, also for the first time, allow for a consistent description of the jets (that have consistent values of the jet flow velocities, radii, opening angles and mass, momentum and energy fluxes) within the cluster heating scenario.  Our simulations show that heating by AGN jets is mainly driven by shock heating, resulting in a very fast and efficient process (more than 95\\% of the energy injected through the jets is transferred to the ambient). As a by-product, our simulations also show that although buoyant cavities could be the very final stage of radio-galactic relics, they are not the main actors in the ICM reheating.  As a result of our simulations, we suggest the idea that most of the observed X-ray cavities are confined by shock waves, very weak though. Although it can not be concluded that the presence of shock confined cavities implies the relativistic nature of jets, such shocks in our results last for much longer periods than in previous -- Newtonian -- works. The confirmation of the existence of such weak shock waves could be an observational challenge that would have crucial implications in our understanding of the galaxy formation and evolution paradigm."
    },
    "1107/1107.5719.txt": {
        "abstract": "{} {We discuss an implementation of our 3D radiative transfer (3DRT) framework with the OpenCL paradigm for general GPU computing. } {We implement the kernel for solving the 3DRT problem in Cartesian coordinates with periodic boundary conditions in the horizontal $(x,y)$ plane, including the construction of the nearest neighbor $\\Lstar$ and the operator splitting step. } {We present the results of a small and a large test case and compare the timing of the 3DRT calculations for serial CPUs and various GPUs. } {The latest available GPUs can lead to significant speedups for both small and large grids compared to serial (single core) computations.} ",
        "introduction": "In a series of papers \\citet*[][hereafter: Papers I--VII]{3drt_paper1, 3drt_paper2, 3drt_paper3, 3drt_paper4, 3drt_paper5,3drt_paper6,3drt_paper7}, we have described a framework for the solution of the radiative transfer equation in 3D systems (3DRT), including a detailed treatment of scattering in continua and lines with a non-local operator splitting method, and its use in the general model atmosphere package \\phx.  The 3DRT framework discussed in the previous papers of this series requires very substantial amounts of computing time due to the complexity of the radiative transfer problem in strongly scattering environments. The standard method to speed up these calculations is to implement parallel algorithms for distributed memory machines using the MPI library \\citep{3drt_paper1,3drt_paper2}. The development of relatively inexpensive graphic processing units (GPUs) with large numbers of cores and the development of relatively easy to use programming models, OpenCL and CUDA has made the use of GPUs attractive for the acceleration of scientific simulations. GPUs are built to handle numerous lightweight parallel threads simultaneously and offer theoretical peak performance far beyond that of current CPUs. However, using them efficiently requires different programming methods and algorithms than those employed on standard CPUs. We describe our first results of implementing our 3DRT framework for a single geometry within the OpenCL \\cite[]{OpenCL} paradigm for generalized GPU and CPU computing. ",
        "conclusions": "We have described first results we have obtained implementing our 3DRT framework in OpenCL for use on GPUs and similar accelerators. The results for 3D radiative transfer in Cartesian coordinates with periodic boundary conditions show that high-end GPUs can results in quite large speed-ups compared to serial CPUs and are thus useful to accelerate complex calculations.  This is in particular useful for clusters where each node has one GPU device and where calculations can be domain-decompositioned with a one node granularity. Large scale calculations that require a domain-decomposition larger than one node are more efficient on large scale supercomputers with 1000's of cores as data transfer required for simultaneous use of multiple GPUs on multiple nodes will dramatically reduce performance.  Even medium-end or low-end GPUs can be useful to offload calculations from the CPU to speed up the overall calculations."
    },
    "1107/1107.2601_arXiv.txt": {
        "abstract": "{High energy emission could be produced in the interaction sites of both galactic and extragalactic jets with the surrounding medium.  We have developed an interaction model that accounts for the continuous injection of relativistic electrons in the forward, reverse and recollimation shocks. We also performed hydrodynamical simulations to establish the physical properties in both type of systems. The resulting non-thermal emission is predicted assuming  different values for the jet power, the external mass density and the source age for both FR-I galaxies and galactic microquasars. The obtained fluxes are compared to current instrument sensitivities at radio, X-ray and gamma-ray bands. We study in detail the connection between hard X-ray synchrotron radiation and gamma-ray emission and use the INTEGRAL, Fermi and AGILE, and Cherenkov telescopes capabilities to test it.}  \\FullConference{8th INTEGRAL Workshop: The Restless Gamma-ray Universe Integral2010,\\\\ September 27-30, 2010\\\\ Dublin Ireland}     \\begin{document} ",
        "introduction": "The jets of Fanaroff-Riley galaxies (of type FR-I and FR-II, \\cite{fr74}) deliver kinetic energy to the surrounding interstellar and intergalactic medium (ISM and IGM respectively) at a rate between $\\sim 10^{42}$ and $\\sim 10^{46}\\,\\rm{erg~s}^{-1}$. The ejections of microquasars ($\\mu$Q), on the other hand, can also inject large amounts of energy into the ISM, at a level $\\sim 10^{37}-10^{39}$~erg~s$^{-1}$. Both of $\\mu$Q and FR-I jet/medium interactions could be strong enough to accelerate particles and produce non-thermal radiation \\cite{br09, br10}. Extended X-ray emission from the jets and/or lobes of the radio galaxies 3C~15 \\cite{ka03}, Cen~A \\cite{cr09}, Fornax~A \\cite{fe95} and M87 \\cite{ka05} have been observed with \\textit{Chandra}. Furthermore, the {\\it Fermi} Collaboration recently reported on the extended GeV emission from Cen~A \\cite{ab10a}.  Efficient particle acceleration is therefore taking place at least in some FR-Is. In the case of $\\mu$Qs, evidence of particle acceleration is found in the non-thermal radio and/or X-ray emission observed e.g. in SS~433 \\cite{zealey80},  XTE~J1550$-$564 \\cite{corbel02} and Cir~X-1 \\cite{tudose06}), and perhaps also in GRS~1915+105 \\cite{Kaiser04} and LS~I~+61~303 \\cite{paredes07}. In this work we explore the leptonic non-thermal emission expected in both $\\mu$Q and FR-I interaction scenarios, using the results from numerical simulations coupled to a radiation model for the jet shocked material (reconfinement region and cocoon) and the ambient material shocked by the bow shock (the shell). Our results are used to make predictions for the broadband non-thermal fluxes (from radio to $\\gamma$-rays) and we compare them  with current and future instrument capabilities. ",
        "conclusions": "The obtained radio fluxes in a $\\mu$Q scenario would imply a flux density of $\\sim 150$~mJy at 5~GHz for a source located at $\\sim$~3 kpc. The emitting size would be of a few arcminutes, since the electron cooling timescale is longer than the source lifetime and they can fill the whole cocoon/shell structures. Considering this angular extension and taking a radio telescope beam size of $10''$, radio emission at a level of $\\sim 1$~mJy~beam$^{-1}$ could be expected. In X-rays, we find a bolometric flux in the range 1--10~keV of $F_{\\rm 1-10 keV}$ $\\sim 2 \\times 10^{-13}$~erg~s$^{-1}$~cm$^{-2}$. The electrons emitting at X-rays by synchrotron have very short time-scales, and the emitter size cannot be significantly larger than the accelerator itself. Although the X-rays produced in the shell through relativistic Bremsstrahlung are expected to be quite diluted, the X-rays from the cocoon would come from a relatively small region close to the reverse shock, and could be detectable by {\\it XMM-Newton} and {\\it Chandra} at scales of few arcseconds. At hard X-rays, \\textit{INTEGRAL} could detect the high energy tail of the particle population, although its moderate angular resolution would make difficult to resolve the sources. At HE, the flux between 100 MeV and 100 GeV is $F_{\\rm 100~MeV<E<100~GeV}\\sim 10^{-14}$~erg~s$^{-1}$~cm$^{-2}$, below the \\textit{Fermi} sensitivity, while the integrated flux above 100~GeV is $F_{\\rm E>100GeV}\\sim 10^{-15}$~erg~s$^{-1}$~cm$^{-2}$, also too low to be detectable by current Cherenkov telescopes. Taking into account the rough linearity between $Q_{\\rm jet}$, $n_{\\rm ISM}$, $t_{\\rm src}$ and $d^{-2}$ with the gamma-ray fluxes obtained, sources with higher values of these quantities than the ones used here may render the $\\mu$Q jet termination regions detectable by current and next future gamma-ray facilities.      For FR-I sources radio fluxes as high as $\\sim 10^{-12}\\,(d/100~$Mpc$)^{-2}$~erg~cm$^{-2}$~s$^{-1}$ or $\\sim 10$~Jy at 5~GHz from a region of few times $10'\\,(d/100~$Mpc$)^{-1}$ angular size are obtained. The properties of the radio emission from the interaction jet-medium structure are comparable with those observed for instance in 3c~15 \\cite{ka03}, in which fluxes of a few $\\times 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ ($d\\sim$~300~Mpc) are found. At X-rays, the dominant emission comes from the cocoon, with fluxes $\\sim 10^{-13}\\,(d/100~$Mpc$)^{-2}$~erg~cm$^{-2}$~s$^{-1}$, although limb brightening effects may increase the shell detectability. The lifetime of X-ray synchrotron electrons, $\\sim 10^{11}\\,{\\rm s}$, is much shorter than in radio, and $\\ll t_{\\rm src}$ as well, so their radiation may come mostly from the inner regions of the cocoon or the bow-shock apex. The total non-thermal X-ray flux is roughly constant for the explored range of $t_{\\rm src}$, although the shell contribution decreases significantly with time. At HE and VHE energies, the obtained fluxes are similar for both the cocoon and the shell, although the latter shows a lower maximum photon energy. Gamma-ray emission increases with time mainly due to the increasing efficiency of the CMB IC channel. The obtained fluxes, for a source with $t_{\\rm src}\\sim 10^8$~yr, are around $\\sim 10^{-12}\\,(d/100\\,{\\rm Mpc})$~erg~cm$^{-2}$~s$^{-1}$. At HE, such a source may require very long exposures to be detected. Nevertheless, {\\it Fermi} has recently detected some FR-I galaxies at a few hundred Mpc distances, including the extended radio lobes of Cen~A (at a distance $\\sim 4$~Mpc), presenting fluxes similar to those predicted here. At VHE, the fluxes would be detectable by the current instruments, although the extension of the source, of tens of arcminutes at 100~Mpc, and the steepness of the spectrum above $\\sim 100$~GeV, may render its detection difficult. In the case of Cen~A, detected by HESS \\cite{ah09}, the emission seems to come only from the core, but this is expected given the large angular size of the lobes of this source, which would dilute its surface brightness, requiring very long observation times to render the source detectable. Finally, we note that the fluxes showed above strongly depend on the non-thermal luminosity fraction going to non-thermal particles, for which we use a quite conservative value $=0.01$. In the case of a source able to accelerate particles at a higher efficiency at the interaction shock fronts, then the expected non-thermal fluxes would be enhanced by a similar factor."
    },
    "1107/1107.5810_arXiv.txt": {
        "abstract": "We show that current microlensing and dynamical observations of the Galaxy permit to set interesting constraints on the Dark Matter local density and profile slope towards the galactic centre. Assuming state-of-the-art models for the distribution of baryons in the Galaxy, we find that the most commonly discussed Dark Matter profiles ({\\it viz.}~Navarro-Frenk-White and Einasto) are consistent with microlensing and dynamical observations, while extreme adiabatically compressed profiles are robustly ruled out. When a baryonic model that also includes a description of the gas is adopted, our analysis provides a determination of the local Dark Matter density, $\\rho_0=0.20-0.56\\textrm{ GeV/cm}^3$ at 1$\\sigma$, that is found to be compatible with estimates in the literature based on different techniques. ",
        "introduction": "\\par The presence of large amounts of Dark Matter (DM) in the Universe, and in particular in the Milky Way, is established on sound observational grounds \\cite{BergstromBook,BertoneReview,BertoneBook}. However, a detailed description of the Dark Matter distribution in the Galaxy is hard to achieve, despite the tremendous progress in numerical simulations of galaxy-sized objects over the last few years \\cite{VL2,VL2site,Aq,Aqsite}, therefore precluding, among other things, a precise interpretation of direct and indirect Dark Matter searches (see e.g.~Refs.~\\cite{StrigariTrotta,Green:2010ri,McCabe:2010zh,paperastrodirect,Vogelsberger,Pieri:2009je}). In order to constrain Milky Way mass models, several observables have been used in the literature, including star counts, the motion of gas and stars or microlensing events. Here, we focus on two key probes, {\\it viz.}~microlensing observations and dynamical measurements, and show that interesting constraints can be set on the DM distribution using those measurements only. This provides complementary evidence for the existence of Dark Matter in the Galaxy, with a distribution compatible with that inferred from other observables.  \\par Gravitational microlensing has long been adopted as a tool to study the structure of the Galaxy \\cite{Paczynski85,KiragaPaczynski}. This technique was suggested in \\cite{Paczynski85} to probe the existence of Massive Compact Halo Objects (MACHOs) by using source stars in the Magellanic Clouds, M31 and M33. Soon after the observation of the first microlensing events in the early 90s \\cite{Alcock:1993eu,Aubourg:1993wb,Udalski:1993zz}, it became apparent that MACHOs could not be a dominant component of the galactic dark halo \\cite{Gates:1994jy} -- therefore a type of dark, non-compact matter was needed. Meanwhile, microlensing of stars towards the galactic centre proved useful in tuning the existent bulge, bar and disk models. Until today, several thousands of microlensing events have been collected -- most notably by OGLE \\cite{Udalski:1993zz,Udalski:1994ei,Sumi:2005da}, MACHO \\cite{Alcock:1993eu,Alcock:2000rv,Popowski:2004uv}, EROS \\cite{Aubourg:1993wb,Afonso:2003aw,Hamadache:2006fw,2009A&A...500.1027R} and MOA \\cite{Sumi:2002wg} campaigns (see \\cite{Moniez:2010zt,CalchiNovati:2009af} for a review) -- allowing for precise estimates of the microlensing optical depth towards the galactic bulge, spiral arms and Magellanic Clouds.  \\par Another well-known tool to constrain the different components of our Galaxy is the rotation curve. Unlike microlensing -- which is sensitive to the distribution of compact matter along the line of sight --, dynamical measurements constrain the total mass distribution. Several decades ago, the observation of flat rotation curves in external spiral galaxies provided convincing evidence for Dark Matter, but actually the rotation curve of the Milky Way (a spiral itself) is far from being precisely determined given our peculiar position within the galactic disk. In fact, although hosts of data have been gathered over the years \\cite{Sofue:2008wt}, allowing for ever more accurate dynamical models, sizeable uncertainties remain on Milky Way mass models, a circumstance that in turn affects our capability to constrain the DM distribution, as we shall see below.  \\par In a pioneering work \\cite{Kuijken}, Kuijken pointed out that a non-zero microlensing optical depth towards the galactic bulge sets a lower limit on the enclosed baryonic mass and thus a lower limit on the corresponding circular velocity within the solar circle. Using this result, Binney \\& Evans \\cite{BinneyEvans} derived an upper limit on the dark matter contribution to the rotation curve, which led them to the conclusion that ``the cuspy haloes favoured by the Cold Dark Matter cosmology (and its variants) are inconsistent with the observational data on the Galaxy'' \\cite{BinneyEvans}. This claim was however based on a preliminary determination of the microlensing optical depth by the MACHO collaboration \\cite{Popowski:2000jq} (later replaced by a less constraining measurement from the same collaboration \\cite{Popowski:2004uv}) as well as a very simplified treatment of the galactic rotation curve.  \\par Here, we revisit this claim in light of recent microlensing observations towards a variety of galactic regions and up-to-date measurements of our Galaxy's rotation curve. A very significant improvement of our work with respect to the existing literature is a proper inclusion of all experimental uncertainties regarding both microlensing and dynamical data, which makes our conclusions sound on statistical grounds. Furthermore, we shall also use a wide range of state-of-the-art models for the galactic baryonic component, so that modelling uncertainties are appropriately bracketed. Our ultimate aim is to combine microlensing and dynamical observables to draw robust constraints on the Dark Matter distribution in the inner region of the Milky Way. ",
        "conclusions": "\\par We have studied the constraints that microlensing and dynamical observations can set on the distribution of Dark Matter in the Galaxy, keeping into account all experimental uncertainties. Starting from state-of-the-art models for the galactic baryonic component, we have rescaled them to match the observed microlensing optical depth towards the galactic bulge, and compared the resulting rotation curve with the one inferred from terminal velocities of gas clouds and other kinematical probes.  \\par This allowed us to revisit the compatibility of different observational probes with the results that emerge from numerical simulations in $\\Lambda$CDM cosmologies. We have followed two different approaches. In the first one, we have set conservative upper limits on the Dark Matter local density and profile shape towards the centre of the Galaxy, working with generalised NFW and Einasto profiles. The fiducial parameters usually adopted in the literature for both profiles have been found to be safely within the allowed regions set by our analysis, contrary to earlier claims of inconsistency between observations and cuspy Dark Matter profiles.  \\par In our second approach, we focussed on the only baryonic model among those discussed here that also contains a description of the amount and distribution of gas, which is expected to provide a non-negligible contribution to the mass and therefore to the rotation curve of the Milky Way. For this specific model, we were able to calculate the values of the Dark Matter parameters that provide the best fit to the combination of microlensing and rotation curve data. The resulting 1$\\sigma$ range for the local DM density, for NFW profiles, was found to be $\\rho_0=0.20-0.56\\textrm{ GeV/cm}^{3}$, therefore consistent with estimates obtained in the literature with different techniques.  \\par Finally, we have studied the consequences of adiabatic compression, often invoked as a mechanism that could increase the amount of Dark Matter in the innermost regions of the Galaxy, and found that although our analysis is not able to discard this mechanism in general, it rules out combinations of local densities and profile slopes invoked in the literature, e.g.~in the context of indirect Dark Matter searches. For NFW profiles with $\\alpha =1.5$ $(1.7)$ we constrain the local DM density to be within the 1$\\sigma$ range $\\rho_0\\simeq0.25-0.34$ $(0.22-0.28)$ GeV/cm$^3$.  \\par As numerical simulations that include the effect of baryons become ever more accurate and reliable, it will be interesting to compare the findings of these simulations with the wealth of observational data that are currently available. In particular, these simulations should allow us to tie together the distribution of baryons and Dark Matter, and to provide a more precise prescription for the slope of the DM profile towards the galactic centre. This slope is currently extrapolated from simulations {\\it without} baryons, which is an unreliable procedure given that baryons are known to dominate the gravitational potential in the inner Galaxy.  \\vspace{0.5cm} \\par {\\it Acknowledgements:} We would like to thank Sebastiano Calchi Novati, Victor P.~Debattista, Juerg Diemand, Silvia Garbari, Esko Gardner, Michael Kuhlen and Justin Read for helpful discussions. F.I.~ is supported from the European Community research program FP7/2007/2013 within the framework of convention \\#235878, and acknowledges the hospitality of the Institute for Theoretical Physics at the University of Z\\\"urich. M.P.~acknowledges the support from Funda\\c{c}\\~{a}o para a Ci\\^encia e Tecnologia (Portuguese Ministry of Science, Technology and Higher Education) under the program POPH co-financed by the European Social Fund in the early stages of this work, and from the Swiss National Science Foundation in the later ones."
    },
    "1107/1107.3577_arXiv.txt": {
        "abstract": "We present an analysis of the mid-infrared (MIR) colours of 165 70$\\mu$m-detected galaxies in the Shapley supercluster core (SSC) at $z{=}0.048$ using panoramic {\\em Spitzer}/MIPS 24 and 70$\\mu$m imaging. While the bulk of galaxies show $f_{70}/f_{24}$ colours typical of local star-forming galaxies, we identify a significant sub-population of 23 70$\\mu$m-excess galaxies, whose MIR colours ($f_{70}/f_{24}{>}25$) are much redder and cannot be reproduced by any of the standard model infrared spectral energy distributions (SEDs). These galaxies are found to be strongly concentrated towards the cores of the five clusters that make up the SSC, and also appear rare among local field galaxies, confirming them as a cluster-specific phenomenon. Their optical spectra and lack of significant ultraviolet emission imply little or no ongoing star formation, while fits to their panchromatic SEDs require the far-infrared emission to come mostly from a diffuse dust component heated by the general interstellar radiation field rather than ongoing star formation. Most of these 70$\\mu$m-excess galaxies are identified as ${\\sim}L^{*}$ S0s with smooth profiles.  We find that almost every cluster galaxy in the process of star-formation quenching is already either an S0 or Sa, while we find no passive galaxies of class Sb or later. Hence the formation of passive early-type galaxies in cluster cores must involve the {\\em prior} morphological transformation of late-type spirals into Sa/S0s, perhaps via pre-processing or the impact of cluster tidal fields, before a {\\em subsequent} quenching of star formation once the lenticular encounters the dense environment of the cluster core. In the cases of many cluster S0s, this phase of star-formation quenching is characterised by an excess of 70$\\mu$m emission, indicating that the cold dust content is declining at a {\\em slower} rate than star formation. We suggest that the excess 70$\\mu$m emission during quenching is due to either: (i) a reduction of the star-formation efficiency as proposed within the morphological quenching scenario; or (ii) a 2--3$\\times$ increase in the dust-to-gas ratio or metallicity of the remaining interstellar medium, as predicted by chemical evolutionary models of galaxies undergoing ram-pressure stripping or starvation. ",
        "introduction": "\\label{intro}  \\setcounter{footnote}{3}  The evolution of a galaxy is a product of both nature and nurture -- both of its mass, and of the environment in which the galaxy finds itself. Understanding the relative importance of nature and nurture remains a key issue in extra-galactic astrophysics. These factors control the star formation rate (SFR) and the dynamical structure or morphology of the galaxy. In isolated galaxies, these observational quantities are largely determined by feedback processes acting within the interstellar medium (ISM) of the galaxy adjoined to its assembly history via mergers. In the harsh cluster environment, star formation activity in infalling spiral galaxies can be both enhanced through collisions \\citep{moss,bretherton} or quenched by various physical mechanisms such as ram-pressure stripping, harassment or starvation which transform them into the passive lenticulars which dominate the dense cluster cores \\citep[for reviews see e.g.][]{treu,boselli,haines07}. These environmental mechanisms produce the well-known star formation (SF)--density \\citep[e.g.][]{dressler85,poggianti99,balogh00,haines09} and morphology--density \\citep{dressler80,dressler97,smith05} relations. However, the dominant evolutionary pathway(s) which generate these relationships remain uncertain.  Moreover, it remains unclear whether these two relations are driven by the same process, by independent processes acting on diverse time-scales, or if one effectively drives the other. While some of the diminution in star-formation seen in cluster galaxies with respect to the field population could be attributed to the different morphological compositions of the two environments, even at fixed morphology, the star-formation rate of cluster galaxies is found to be lower than in the field \\citep{balogh00}, indicating separate processes at least partially shape these relations.  Alternatively, \\citet{martig} suggest that it is the morphological transformation of spirals into S0s and the growth of the stellar bulge which drives the quenching of star-formation in lenticulars by making the gas disk more stable against collapse and fragmentation into star-forming clumps, without requiring gas consumption or removal. This interpretation is supported by the trend for the molecular gas consumption time-scale to increase systematically with galaxy concentration index ($R_{90}/R_{50}$), recently observed by \\citet{saintonge}. This quenching itself is not related to the environment of the galaxy, and so cannot alone explain the SFR--density relation, instead requiring some additional environmental process(es) capable of producing the prior morphological transformation of the galaxies, such as the impact of cluster tidal fields \\citep{byrd}, or pre-processing within infalling groups \\citep{balogh04b}.  Infalling spirals in local clusters show evidence for the ongoing removal of gas via ram-pressure stripping \\citep{crowl}, leading to a rapid termination of star-formation from the outside-in, and subsequent fading of the disk component. However, S0s are found to differ from normal spirals due to higher bulge luminosities rather than fainter disks \\citep{christlein}, requiring bulge growth during S0 formation, disfavouring ram-pressure stripping or starvation models. Instead the increased scatters seen in cluster S0 and spiral Tully-Fisher relations support mechanisms that kinematically disturb infalling spirals, such as merging or harassment \\citep{moran07}. Such tidal interactions are capable of channelling material to a central bulge sufficient to produce the higher central mass densities ($V_{rot}^{2}/Gr_{e}$) seen in cluster spirals \\citep{moran07} and ultimately the central surface brightnesses and stellar phase densities found in S0s.  A number of studies have attempted to distinguish among the various environmental processes by identifying and classifying various promising classes of transition galaxies, such as post-starburst and/or E+A galaxies \\citep{dressler83,poggianti99,mercurio04,mahajan,mahajan11}, interaction-induced starbursts \\citep{moss,fadda}, and passive/red spirals \\citep{vandenbergh,bamford}, indicating that more than one mechanism is required.  A key diagnostic to understand the relevance of these proposed transition galaxies is mid-infrared (MIR) data to account for the dust-obscured star formation. Many ``post-starburst'' galaxies identified by their deep H$\\delta$ absorption lines and lack of O{\\sc ii} emission were revealed by mid-infrared data to be dusty starbursts \\citep{duc,dressler09}, whose age-dependant dust obscuration preferentially hides the young stars responsible for the O{\\sc ii} emission above those older stars responsible for the Balmer absorption \\citep{poggiantiwu}. These spectroscopic and MIR surveys revealed a ubiquitous population of starburst and post-starburst ${\\sim}L^{*}$ spirals in $z{\\sim}0.5$ clusters \\citep{couch,geach,dressler09}, which are absent in present epoch clusters \\citep{poggianti99}. They have been identified as the likely progenitors of present day cluster S0s, whose populations have risen dramatically over the last 5\\,Gyr, empirically replacing the spirals \\citep{dressler97,treu,smith05}. Additionally, interaction-induced starbursts are often highly obscured \\citep{mihos}, while many red sequence galaxies (including spirals) have been revealed to be actively star-forming, their red colours being attributed to dust obscuration rather than being passively-evolving \\citep{wolf05,wolf,haines10}.  Most previous studies of environmental effects on galaxies have concentrated on their impact on morphology and star formation. However, given that the environmental processes behind the SF--density relation are likely to act {\\em directly} on the ISM rather than the star forming regions themselves, to fully understand the origin of the SF-density relation requires also measures of these impacts of environment on the ISM which subsequently drive it. To date, a detailed knowledge of the impact of environment on the ISM and dust content, and how this drives the SFR--density relation is still largely lacking.  The most accessible probe of the ISM in distant galaxies is via far-infrared observations of the thermal emission from their dust contents. Using early {\\em IRAS} \\citep{iras} and sub-mm observations \\citet{desert} modelled the origin of the far-infrared emission in terms of two main components (ignoring PAHs here): a warm component of small dust grains in H{\\sc ii} regions heated by star formation; and a cool cirrus component of large dust grains heated by the interstellar radiation field. \\citet{lonsdale} estimated that this cool, cirrus component contributes as much as 50--70 per cent of the far-IR flux, while \\citet{sauvage} revealed a strong morphological trend, with a decreasing cirrus contribution from ${\\sim}85$ per cent for Sa galaxies to just 3 per cent for Sdm galaxies. The unprecedented wavelength coverage and sensitivity of {\\em Herschel} \\citep{pilbratt} allows us to extend these analyses to large samples of galaxies. \\citet{kramer} demonstrate the requirement of a two-component dust model to fit the 24--500$\\mu$m spectral energy distribution (SED) of M33, comprising a hot component (${\\sim}6$0K) tracing star-formation and a dominant cold component of large dust grains heated by the general interstellar radiation field. This second component can in effect measure the fuel available for star formation, and a comparison of the two components allows us to measure the star-formation efficiency, or gas consumption time-scale.  Using {\\em IRAS}, \\citet{doyon} found that H{\\sc i}-deficient galaxies in the Virgo cluster had lower 60 and 100$\\mu$m fluxes, and cooler far-infrared colours than those with normal H{\\sc i} content. They interpreted this as evidence for the same two-component model for the far-infrared emission. In the cluster environment, the warm star-formation component is diminished in proportion to the H{\\sc i} deficiency suggesting that star-formation is regulated by the amount of available gas, while the cool component from diffuse dust is reduced by a smaller factor as dust is stripped from the galaxy. {\\em ISO} extended the far-infrared coverage beyond that of {\\em IRAS} to 240$\\mu$m, revealing the presence of a significant cold dust component (15--20K) in a complete sample of 38 late-type Virgo cluster galaxies \\citep{popescu02}.  While the mid-infrared coverage of {\\em Spitzer}/MIPS is dominated by the reprocessed dust emission from current star formation, particularly at 24$\\mu$m, for the 70$\\mu$m band typically ${\\sim}4$0\\,per cent of the integrated light from galaxies can be attributed to dust heated by the diffuse interstellar radiation from evolved stars \\citep{li10,bendo,calzetti10}. Hence the mid-infrared colours ($f_{70}/f_{24}$) can be considered as a sensitive probe of the interrelation between the ISM and star formation in galaxies.   In this paper we examine the mid-infrared colours of galaxies in the Shapley supercluster core (SSC) at $z{=}0.048$, the most massive and dynamically active structure in the local ($z{<}0.1$) Universe \\citep{raychaudhury}. In this paper the SSC refers to a chain of three merging clusters, A3558, SC\\,1327-312 and SC\\,1329-313, along with the two associated Abell clusters on either side, A3562 and A3556. In {\\S}~2 we present the panoramic {\\em Spitzer} and ancillary data obtained within the ACCESS\\footnote{{\\em http://www.na.astro.it/ACCESS}} multi-wavelength survey used in this work. In {\\S}~3 we present the mid-infrared colours ($f_{70}/f_{24}$) of galaxies in the SSC, identifying a subpopulation with much redder infrared colours ($f_{70}/f_{24}{\\ga}25$) than can be reproduced by any of the standard infrared SED models, before examining in detail the nature of these objects. In {\\S}~4 we discuss possible origins of these 70$\\mu$m-excess galaxies, which appear to be dust-rich S0s yet with extremely low SFRs, before examining their consequences for the possible evolutionary pathways in the formation of cluster S0s. When necessary we assume \\mbox{$\\Omega_{M}{=}0.3$}, \\mbox{$\\Omega_{\\Lambda}{=}0.7$} and \\mbox{H$_{0}{=}70\\,$km\\,s$^{-1}$Mpc$^{-1}$}, such that at the distance of the Shapley supercluster 1\\,arcsec is equivalent to 0.96\\,kpc. ",
        "conclusions": "\\label{sec:discuss}  We examined the mid-infrared colours of 165 70$\\mu$m-detected spectroscopic members of the Shapley supercluster core at $z{=}0.048$, using panoramic {\\em Spitzer}/MIPS 24$\\mu$m and 70$\\mu$m data. Among these galaxies we identified a population of 23 ``70$\\mu$m-excess'' galaxies, comprising $14{\\pm}3$ per cent of 70$\\mu$m-detected SSC galaxies, whose mid-infrared colours $(f_{70}/f_{24}{>}25)$ are much redder than can be reproduced by any of the standard infrared SED models, and which also appear extremely rare in field surveys. These galaxies are primarily massive ${\\sim}L^{*}_{K}$ galaxies with E/S0 morphologies, and intermediate $f_{24}/f_{K}$ ratios placing them in transition zone in which galaxies in the process of having their star formation quenched are likely to be located. In the two-component model for the far-infrared emission these colours could be understood as galaxies in which the star-formation is being quenched (or was recently), thus reducing the 24$\\mu$m emission and explaining the low $f_{24}/f_{K}$ flux-ratios, while the emission seen at 70$\\mu$m is largely due to dust heated predominately by relatively young stars ($10^{8}$--$10^{9}$\\,yr) and the interstellar radiation field.  These 70$\\mu$m-excess galaxies comprise $8{\\pm}2$ per cent of all $K{<}13$ ($M_{K}{<}M^{*}{+}1.3$) SSC members, indicating that they represent an important evolutionary pathway for cluster galaxies. Given this, we now attempt to place these galaxies within the theoretical frameworks provided by the various environmental processes that are currently believed to be important in transforming galaxies within clusters, including ram-pressure stripping and morphological quenching.  \\subsection{Ram-pressure stripping}  The concentration of these 70$\\mu$m-excess galaxies in the cluster cores would suggest ram-pressure stripping as the mechanism causing the observed quenching of star-formation. In such galaxies, the gas is stripped from the outside-in, leaving truncated H{\\sc i} disks \\citep{vogt}. Within this truncation radius star formation continues normally, but outside evidence is found for recently quenched stellar populations \\citep{crowl}. Using spatially-resolved {\\em Herschel}/SPIRE maps \\citet{cortese10} show the cold dust to be collocated with the gas, appearing truncated for highly H{\\sc i}-deficient spirals in the Virgo cluster. Hence, while star formation is quenched and gas removed via ram-pressure stripping, as the gas/dust surface densities within the truncation radius remain high, the star-formation {\\em efficiency} should be unaffected, and the SFR and gas contents should decline in step.  While the H{\\sc i} gas contents of cluster galaxies may be partially or completely stripped by ram pressure, the molecular gas, by virtue of its much higher density and being more tightly bound within the potential well of the galaxy, should be much more difficult to remove. Indeed, in a CO survey 260 early-type galaxies from the volume-limited ATLAS$^{\\rm 3D}$ sample, \\citet{young} recently found no significant difference between the molecular gas contents of early-type galaxies in the Virgo cluster and those in isolated field regions within 24\\,Mpc. Moreover, the dynamical status of the 12 CO-detected Virgo cluster galaxies (all S0s) is consistent with being a virialized population, deeply embedded within the hot intracluster gas, and unlikely to be recently accreted cluster members. They have probably retained their molecular gas over several Gyr within the cluster, albeit their molecular gas fractions (M$_{{\\rm H}_{2}}/{\\rm M}_{*}{\\sim}0.$001--0.01) are now a factor ${\\ga}10{\\times}$ lower than those of normal star-forming galaxies \\citep{saintonge}.  As ram-pressure stripping and the milder starvation mechanism prevent further accretion of pristine gas, the remaining H{\\sc i} gas is enriched by metals recycled from stellar mass loss, particularly via the shedding of the stellar envelopes during the asymptotic giant branch phase, increasing the gas metallicity by a factor $\\sim$2--3 \\citep{skillman,boselli08}, which may in turn reduce the star formation efficiency (and hence 24$\\mu$m emission) by a similar factor \\citep{dib}. As the dust-to-gas ratio is known to scale approximately linearly with metallicity \\citep{draine07}, we may expect the ratio between dust content and SFR to increase by a similar factor 2--3. This however requires us to make many assumptions about how the metals produce by evolved stars are frozen into dust grains as well as the rate at which dust grains are destroyed in the ISM for S0s passing through the cores of clusters as opposed to isolated field spirals. This enrichment of the remaining ISM by stellar mass loss could though be responsible for the anomalously red $f_{70}/f_{24}$ colours. Moreover an increase in the dust-to-gas ratio may directly reduce the star formation efficiency, if radiation pressure from the absorption and scattering of starlight by dust grains acts as an important mechanism in regulating star formation \\citep{andrews}.  \\citet{leitner} show that at low-redshifts gas recycling from stellar mass loss can provide most if not all of the gas required for star formation in galaxies, particularly those with lower than average specific-SFRs (as our galaxies are). The key for stellar mass loss to be efficient in fuelling star formation is the presence of at least some gaseous disk, into which the enriched material can be released directly. For elliptical galaxies with hot gas halos, the enriched gas released by AGB stars is likely rapidly disrupted and mixed with the hot halo gas. The simulations of  \\citet{martig10} suggest that continuous recycling of gas through stellar mass loss also contributes significantly to the growth and survival of the disk component in these cluster S0s. \\citet{davis} find that the molecular gas of the Virgo cluster CO-detected S0s is always kinematically aligned with the stellar kinematics, consistent with an internal origin such as stellar mass loss.  \\begin{figure} \\centerline{\\includegraphics[width=80mm]{rps.eps}} \\caption{Effective ram pressure ($\\rho_{ICM}\\nu_{orb}^{2}$) acting on galaxies infalling into Abell 3558 as a function of cluster-centric radius (solid line), based on the gas/mass density profiles estimated by \\citet{sanderson10} from {\\em Chandra} X-ray data. The shaded regions indicate the 1$\\sigma$ errors extrapolated from the uncertainties in the \\citet{ascasibar} model parameters obtained from 200 bootstrap resamplings of the original X-ray data \\citep[for details see][]{sanderson10}.} \\label{rps} \\end{figure}   Another explanation could be that the cold dust is heated during the stripping process, boosting the 70\\,micron fluxes. This would be consistent with the 100$\\mu$m-excess galaxies found by \\citet{rawle} in the Bullet cluster, as the dust temperatures estimated by their fitting their combined {\\em Herschel}/PACS+SPIRE photometry for these galaxies were found to be higher than expected given their infrared luminosity. The ram-pressure could also be removing the dust/gas from the vicinity of the source of heating (the stars). \\citet{cortese10b} found extra-planar cold dust in the perturbed galaxy NGC\\,4438 in Virgo collocated with the stripped atomic and molecular hydrogen. This extra-planar dust lacks a warm component or associated star formation, leading to unusual $f_{70}/f_{24}$ colours similar to those presented by our 70$\\mu$m-excess galaxies.  Finally, we check to see if ram-pressure stripping should be effective for these galaxies. For Abell 3558, \\citet{sanderson10} derived the dark matter density, gas density ($\\rho_{ICM}$) and gas temperature profiles by fitting \\citet{ascasibar} models to the {\\em Chandra} X-ray data. In these models, the dark matter is described by a \\citet{hernquist} density profile, which provides simple analytic expressions for the mass profile, gravitational potential and escape velocity ($\\nu_{esc}$). From this, we can derive the ram-pressure, $P_{ram}{=}\\rho_{ICM}\\nu_{orb}^2$ acting on galaxies as a function of cluster-centric radius, as shown in Fig.~\\ref{rps}. For the orbital velocities, we assume the galaxies start at rest at the turnaround radius ($r_{ta}{\\sim}3.5\\,r_{200}$) and fall radially. We find obtain values of ${\\sim}1\\,000\\,{\\rm cm}^{-3}({\\rm km\\,s}^{-1})^{2}$ at 0.68\\,$r_{500}$ \\citep[where $r_{500}{=}1.29$\\,Mpc;][]{sanderson}, indicative of moderate ram-pressure stripping, capable of partially stripping an $L^{*}$ spiral with rotational velocity $\\nu_{rot}{\\sim}1$50--200\\,km\\,s$^{-1}$ \\citep[e.g.][]{vollmer,roediger,kronberger}, rising to ram pressures of ${\\sim}10\\,000\\,{\\rm cm}^{-3}({\\rm km\\,s}^{-1})^{2}$ at 0.24\\,$r_{500}$, at which point all of the gas may be rapidly stripped from such galaxies. Thus for the bulk of our 70$\\mu$m-excess galaxies, which lie within ${\\sim}0$.5--1.0\\,$r_{500}$, moderate-to-strong ram-pressure stripping should be effective.  Furthermore, the velocity dispersion of the 13 70$\\mu$m-excess galaxies within $r_{500}$ of Abell 3558 is 948\\,km\\,s$^{-1}$, marginally lower than the $975{\\pm}39$\\,km\\,s$^{-1}$ obtained for the global cluster population within $r_{500}$. This would indicate that these 70$\\mu$m-excess galaxies are a virialized population, having been embedded in the dense ICM for a Gyr or more, rather than being on their first infall into the cluster. The fact that they make up $8{\\pm}2$ per cent of the cluster galaxy population would also imply that this 70$\\mu$m-excess phase is long-lived, rather than the rapid evolution typically expected from ram-pressure stripping of infalling spirals, but as \\citet{young} argue, the molecular gas content should be much more resistant to ram-pressure stripping than the easily removed H{\\sc i} component.  \\subsection{Morphological quenching}  The 70$\\mu$m-excess galaxies are morphologically identified as S0s, with smooth profiles and absent spiral arms. An alternative mechanism that could cause the decline in star-formation to occur more rapidly than the reduction in the dust/gas content could be morphological quenching. In this hypothesis, recently developed by \\citet{martig}, the morphological transformation from spiral to lenticular makes the gas (and collocated dust) more stable against fragmentation and collapse into molecular clouds, leading to a quenching of star formation for a constant gas/dust fraction. The 24$\\mu$m flux tracing the star-formation continues to decline, while the 70$\\mu$m flux comes predominately from the cold interstellar dust content. The general interstellar radiation field in the bulge is predicted to be sufficiently strong to heat this dust without requiring young stars due to the high central density of stars \\citep{jura}. IRAS detected far-infrared emission from a large fraction of nearby early-type galaxies, showing that they contained diffuse cool interstellar matter \\citep{knapp,goudfrooij,merluzzi98}. Recent {\\em Herschel} observations of nearby galaxies by \\citet{engelbracht} and \\citet{bendo} show that the temperatures of the cool dust component in the central regions of early-type spirals (S0--Sb) are 20--50\\% higher than in the disks, indicating that the dust is heated by the higher stellar density in the galactic bulge. Again this dust heating would be consistent with the SED fits of \\citet{rawle} to their 100$\\mu$m-excess cluster population, although only one of their five 100$\\mu$m-excess galaxies with HST imaging appears bulge-dominated.  While early-type galaxies are classically defined as gas-poor passive systems, molecular gas, H{\\sc i} and warm dust was detected in 30 per cent of the SAURON representative sample of E/S0 galaxies \\citep{combes}. \\citet{crocker} resolved the molecular gas for 12 SAURON galaxies into central disks or rings, coincident with the dust distribution (often in the form of tightly-wound spiral structures) and corrotating with the ionized gas. However, star formation is the ionisation source for just half of the sample, corresponding to those with the largest molecular gas reservoirs. Most of their E/S0s do not lie on the FIR--radio relation, having excess FIR flux, while the bolometric IR flux produces SFR estimates ${\\sim}3$ times higher than those based on the 8$\\mu$m or 24$\\mu$m fluxes, due to additional heating from evolved stars in the bulge. Interestingly, one galaxy from their sample, NGC\\,4526 an S0 belonging to the Virgo cluster, is also a 70$\\mu$m-excess galaxy with $f_{70}/f_{24}{\\sim}30$, supporting the hypothesis that these are a cluster population. Of the 225 mostly field E/S0 galaxies in the {\\em Spitzer}/MIPS study of \\citet{temi09} only NGC\\,5866 and NGC\\,4526 have $f_{70}/f_{24}{>}30$, again suggesting that this is a rare or short-lived phenomenon in isolated field E/S0s. The above results all reveal that a fraction of early-type galaxies have significant gas and dust contents, producing emission in the far-infrared, at levels above those predicted given their extremely low levels of star formation.  \\begin{figure} \\centerline{\\includegraphics[width=80mm]{optmir_morphology.eps}} \\caption{The infrared colour $f_{24}/f_{K}$ versus $K$-band magnitude for confirmed Shapley supercluster galaxies, colour-coded according to their morphologiesas indicated. The red dashed lines indicate how we separate galaxies into star-forming, transition and passive populations according the their $f_{24}/f_{K}$ ratio as discussed in the text. Those galaxies not detected at 24$\\mu$m are also plotted, after being given a random 24$\\mu$m flux in the range 100--200$\\mu$Jy, that is a factor 2--4 below the survey limit, and appear along the diagonal sequence passing through $K{\\sim}13.5$, $f_{24}/f_{K}{\\sim}0.05$.} \\label{optmir_morph} \\end{figure}   \\subsection{Caveats with both scenarios}  The presence of numerous S0s within the star-forming sequence of Fig.~\\ref{optmir} and all along the quenching sequence of Fig.~\\ref{quenching}, appears to point towards a morphological scenario in which spirals are first transformed into S0s without necessarily affecting their star-formation, before being quenched morphologically. This however, doesn't explain why these transition dusty S0s with litte or no star formation are located almost exclusively in the dense cluster cores and appear so rare in isolated field regions. In contrast the ram-pressure stripping scenario is able naturally to explain their concentration in the cluster cores where the dense ICM and high infall velocities make ram-pressure stripping most efficient. It doesn't however explain why these galaxies have already been morphologically transformed from spirals to S0s, as ram-pressure stripping is not predicted to strongly affect a galaxy's morphology. Instead, we would expect to find large numbers of spirals with little or no star-formation, which judging by Fig.~\\ref{quenching} isn't the case. However we note that galaxies for which star-formation has been largely or completely quenched would not likely be detected at 70$\\mu$m and hence aren't included in the figure.  \\subsection{Relating the quenching of star formation to the morphological transformation of galaxies within clusters}  To accurately account for the quenched galaxy population, and to be able to determine if indeed star-formation is primarily quenched {\\em after} morphological transformation rather than before, we show in Fig.~\\ref{optmir_morph} the infrared colour $f_{24}/f_{K}$ versus $K$-band magnitude of confirmed Shapley supercluster galaxies, colour-coded according to their morphological classifications, allowing us to directly relate the bimodality in the star-formation activity to morphology. As we would expect, the passive galaxies ($f_{24}/f_{K}{<}0.15$) are predominately early-types (E,E/S0,S0), while the normal star-forming galaxies ($f_{24}/f_{K}{>}1.0$) are mostly spirals (Sa or later). There are however notable populations of both passive spirals and star-forming S0s.  Looking first at the 277 passive galaxies with $K{<}14$, we find that $10.1_{-1.5}^{+2.1}$ per cent are classified as spirals (27 Sa, 1 Sg). These examples of massive spirals with low SFRs are reminiscent of the classic ``anemic spiral'' of \\citet{vandenbergh}, although all show also a prominent bulge. This would be consistent with \\citet{bell} who show that truly quenched central galaxies (i.e. most massive galaxies within their host DM halo) have stellar bulges, indicating that a bulge, is an absolute requirement for full quenching of star formation (at least in central galaxies). For these passive spirals certainly ram-pressure stripping is a suitable candidate for removing their gas and quenching star-formation, in a manner similar to those spirals with truncated H$\\alpha$-disks and/or H{\\sc i} tails found in Virgo \\citep{vogt,koopmann,crowl,chung}. Conversely, considering the 77 $K{<}14$ star-forming galaxies, $18.2_{-3.6}^{+5.2}$ per cent are classified as early types (2 S0/S, 10 S0, 1 E/S0, 1 E), for which it seems the morphological transformation has occured {\\em prior} to any significant quenching of star-formation. This time sequence is the opposite of what has been suggested by previous studies \\citep[e.g.][]{balogh00,bamford} which typically use optical colour as a proxy for star-formation activity. However, using {\\em Spitzer} mid-infrared data \\citet{wolf} revealed that these optically passive spirals still had significant obscured star formation, albeit at levels 0.6\\,dex lower on average than normal star-forming galaxies of the same mass. They suggested that these apparent contraditions between the optical and mid-infrared viewpoints reflected that the unobscured star-formation in the outer disk is the most susceptible to quenching by ram-pressure stripping, while the obscured component in the central parts remains relatively unaffected.  That there are comparable fractions and numbers of passive spirals and star-forming early-types indicates that star-formation in cluster galaxies is being quenched {\\em both} before {\\em and} after morphological transformation from spiral to S0. This is confirmed if we look at the morphologies of the transition galaxies with $0.15{<}(f_{24}/f_{K}){<}1.0$ likely to be in the process of having their star formation quenched. Here we find that $41.9_{-5.5}^{+5.9}$ per cent of the 74 $K{<}14$ transition galaxies are spirals (26 Sa, 2 Sb, 3 Sg), and $58.1_{-5.9}^{+5.5}$ per cent are early-types (29 S0, 2 S0-S, 6 E/S0, 6 E). Among this latter population, many have excess 70$\\mu$m emission, producing mid-infrared colours inconsistent with those seen in the local field population or by model infrared SEDs.  The concentration of these transitional S0s with excess 70$\\mu$m emission towards the cluster cores points towards ram-pressure stripping being as the mechanism quenching the star formation in these galaxies. However, assuming the simple \\citet{gunn} criterion the high central stellar densities of these S0s should significantly reduce the ability of ram-pressure stripping to completely remove their gas, while at low redshifts stellar mass loss should be capable of replenishing the gas supply requried to maintain observed low levels of star formation. As the findings of \\citet{young} show, a significant fraction of S0s are able to retain their molecular gas for several Gyr within the dense ICM of the Virgo cluster. The apparent ``virialized'' dynamical status of the 70$\\mu$m-excess population, suggests that these have also resided within the ICM for some time, rather than observed on their first infall. If this is the case, then it seems plausible that these are the direct descendants of the dusty starburst and post-starburst ${\\sim}L^{*}$ galaxies abundant in clusters at $z{\\sim}0.4$ \\citep{poggianti99,dressler09}, but which appear rare in present day clusters."
    },
    "1107/1107.4357_arXiv.txt": {
        "abstract": "We study the frequency of \\mgii absorption in the outer haloes of galaxies at $z=0.6-1.4$ (with median $z=0.87$), using new spectra obtained of ten background quasars with galaxy impact parameters of $b<100$ kpc.  The quasar sightlines were selected from the SDSS DR6 QSO catalog based on proximity to galaxies in the DEEP2 redshift survey.  In addition to the 10 small impact systems, we examine 40 additional galaxies at $100<b<500$ kpc serendipitously located in the same fields.  We detect \\mgii absorbers with equivalent width $W_r$ = 0.15 \\AA\\ - 1.0\\AA, though not all absorbers correlate with DEEP galaxies.  We find five unique absorbers within $\\Delta v = 500$ km/s and $b<100$ kpc of a DEEP galaxy; this small sample contains both early and late type galaxies and has no obvious trends with star formation rate.  No \\mgii is detected more than 100 kpc from galaxies; inside this radius the covering fraction scales with impact parameter and galaxy luminosity in very similar fashion to samples studied at lower redshift.  In all but one case, when \\mgii is detected without a spectroscopically confirmed galaxy, there exists a plausible photometric candidate which was excluded because of slit collision or apparent magnitude.  We do not detect any strong absorbers with $W_r>1.0$\\AA, consistent with other samples of galaxy-selected \\mgii systems.  We speculate that \\mgii systems with $0.3<W_r<1.0$ trace old relic material from galactic outflows and/or the halo assembly process, and that in contrast, systems with large $W_r$ are more likely to reflect the more recent star forming history of their associated galaxies. ",
        "introduction": "For roughly two decades, a connection has been observed between $\\sim L^*$ galaxies and the \\mgii absorption systems seen in QSO spectra \\citep{bergeron_boisse, steidel,steidelandsargent}.  \\mgii systems trace gas in the outer haloes of these objects out to radii of $\\sim 100$ kpc.  However, a standardized physical picture describing the origin of this metal-rich material remains elusive.  At least two different pathways have been suggested for distributing \\mgii throughout galaxy haloes.  In one scenario, the QSO absorbers represent debris from large-scale winds that are seen in galaxy spectra and known to carry \\mgii \\citep{ubiquitous2009, bond}.  These winds may help quench star formation in nascent galaxies, providing necessary feedback in galaxy formation models.  For example, \\citet{nestor2010strong} associated ultrastrong absorbers ($W_r \\ge$ 3\\AA) with starbursting galaxies and suggested an outflow wind scenario based on velocity dispersions in the \\mgii systems.  \\citet{prochtersdss2006} favored an outflow model of \\mgii systems based on evolutionary trends with redshift in the number density of (strong) absorbers.  Some theories associate these outflows with AGN activity rather than supernova feedback, but differentiation between the two remains an open question.  Several authors have argued indirectly in favor of this interpretation as well.  \\citet{bouche2006anticorrelation} detected evidence that stronger \\mgii absorbers reside in lower mass galaxy haloes, and suggested that this might result from the more vigorous star formation occurring in the smaller, late-type galaxies.  Using various statistical combinations of SDSS \\mgii absorbers, \\citet{zibetti2007stacking} and \\citet{menardo2} reported evidence of star formation surrounding the strongest \\mgii systems, in the form of enhanced [OII] emission or bluer galaxy colors.  However, a second interpretation posits that \\mgii systems are a generic feature of galaxy haloes, with an extent and dynamical structure that reflects gravitational processes.  Indeed the velocity profiles of some \\mgii systems appear to form an extension of galactic rotation curves \\citep{kacprzak2010sim, mgii_rotation}, albeit with increased velocity dispersion.  \\citet{chenandtinker2008} and \\citet{chen2010lowz} develop a halo-motivated model using a large set of QSO-galaxy pairs.  These papers and \\citet{gauthier} interpret the anti-correlation between galaxy halo mass and \\mgii strength in the context of the ``cold mode'' vs. ``hot mode'' accretion models recently advocated by \\citet{keres, dekel}.  In this model, \\mgii absorption arises in cold structures that are ionized away when galaxy masses - and hence virial temperatures - exceed a threshold value.  In a carefully selected sample, \\citet{chenandtinker2008} find absorbers associated with both early- and late-type galaxies in normal field proportions, and use this to argue against an interpretation where \\mgii arises exclusively in starburst galaxies.  Galaxy-absorber studies to date have primarily focused on low redshifts $z\\lesssim0.5$ where galaxy observations are most accessible \\citep{chen2010lowz}, or on targeted galaxy observations toward sightlines pre-selected for \\mgii absorption \\citep{kacprzak2010sim, nestor2005, mgii_rotation}.  A uniform interpretation of the results is also complicated by the different samples' varying sensitivities, completeness estimates, and selection techniques.  In this paper, we construct a modest sample for studying \\mgii-galaxy associations at higher redshift ($0.6<z<1.4$).  To do this, we leverage the public catalog of the DEEP2 redshift survey \\citep{deep2_1,deep2_2} and use the MagE spectrometer on Magellan to obtain spectra of background QSOs that lie randomly near DEEP2 galaxies on the sky.  Since this galaxy sample is chosen without regard to morphology, it is not biased toward any spectral type or the presence/absence of AGN; thus, it should be appropriate for investigating the relationship between \\mgii incidence, star formation, and spectral/morphological type.  The sample also spans a wide range of impact parameters, allowing for an estimate of how the \\mgii covering fraction scales with galactic radius.  The DEEP sample does pre-select galaxies based on color and magnitude and is known to preferentially select star forming galaxies at high redshift \\citep{gerke}.  However, it provides a large, uniformly selected parent sample against which we can compare a subsample drawn randomly based only on impact parameter.  In Section 2, we describe our sample selection and observations. Section 3 details the galaxy-\\mgii pairs detected as well as non-detections.  Section 4 describes our results in the context of the existing literature.  We use a WMAP7+BAO+$H_0$ cosmology of $\\Omega_m = 0.272, \\Omega_{\\Lambda} = 0.728, H_0 = 70.4$ km s$^{-1}$ Mpc$^{-1}$ throughout.  All magnitudes are on the AB system, and impact parameters are listed in proper $h_{71}^{-1}$ kpc. ",
        "conclusions": ""
    },
    "1107/1107.4343_arXiv.txt": {
        "abstract": "The Arecibo L-Band Feed Array Zone of Avoidance (ALFA ZOA) Survey has discovered a nearby galaxy, ALFA ZOA J1952+1428, at a heliocentric velocity of +279 $\\mathrm{km \\; s^{-1}}$. The galaxy was discovered at low Galactic latitude by 21-cm emission from neutral hydrogen (H\\,{\\sc{i}}). We have obtained follow-up observations with the EVLA and the 0.9-m SARA optical telescope. The H\\,{\\sc{i}} distribution overlaps an uncataloged, potential optical counterpart. The H\\,{\\sc{i}} linear size is 1.4 kpc at our adopted distance of D = 7 Mpc, but the distance estimate is uncertain as Hubble's law is unreliable at low recessional velocities. The optical counterpart has m$_B$ = 16.9 mag and $B$ - $R$ = 0.1 mag. These characteristics, including M$_{HI}$ = 10$^{7.0}$ M$_\\odot$ and $L_{B}$ = $10^{7.5}$ L$_\\odot$, if at 7 Mpc, indicate that this galaxy is a blue compact dwarf, but this remains uncertain until further follow-up observations are complete. Optical follow-up observations are ongoing and near infrared follow-up observations have been scheduled. ",
        "introduction": "The Arecibo L-Band Feed Array Zone of Avoidance (ALFA ZOA) Survey searches for 21-cm line emission from neutral hydrogen (H\\,{\\sc{i}}) in galaxies behind the disk of the Milky Way. The survey  uses the ALFA receiver on the 305-m Arecibo Radio Telescope\\footnote{Arecibo Observatory is part of the National Astronomy and Ionosphere Center, which, when the observations were made, was operated by Cornell University under a cooperative agreement with the NSF.}. This region of the sky is termed the Zone of Avoidance by extragalactic astronomers because of its low galaxy detection rate. Extragalactic observations at visual wavelengths struggle with high extinction levels. Near and far infrared observations suffer confusion with Galactic stars, dust, and gas. 21-cm line observations are sensitive to late-type galaxies in general and are not affected by extinction. As a spectral line survey, we generally only have confusion with Galactic H\\,{\\sc{i}} within approximately $\\pm$100 $\\mathrm{km \\; s^{-1}}$. The ALFA ZOA survey is sensitive to galaxies behind the Milky Way that go undetected at other wavelengths. It has been suggested by Loeb and Narayan (2008) that undiscovered mass behind the Milky Way may explain the discrepancy between the cosmic microwave background dipole and what is expected from the gravitational acceleration imparted on the Local Group by matter in the local universe (Erdogdu et al. 2006).  Two large area H\\,{\\sc{i}} ZOA surveys have preceded ALFA ZOA; the Dwingeloo Obscured Galaxies Survey and the HI Parkes Zone of Avoidance Survey (HIZOA). The Dwingeloo survey detected 43 galaxies in the northern hemisphere within $\\pm 5^\\circ$ of the Galactic plane. It was sensitive only to nearby, massive objects because of its relatively high noise level of 40 mJy beam$^{-1}$ (with velocity resolution of 4 km s$^{-1}$; Henning et al. 1998). More recently, HIZOA covered decl. = -90 to +25 at 6 mJy/beam rms (with velocity resolution of 27 km/s), and detected about 1000 galaxies (Donley et al. 2005; Henning et al. 2000, 2005, Shafi 2008.)  The ALFA ZOA survey is being conducted in two phases: a shallow and a deep phase. The shallow phase (rms = 5 mJy with velocity resolution of 10 km/s) covers 900 square degrees through the inner Galaxy ($30^\\circ<l<75^\\circ$, $|b|<10^\\circ$) and is expected to detect 500 galaxies. Hundreds of galaxies have been detected so far, and data reduction and analysis are ongoing. This is complemented by a deep survey ($30^\\circ<l<75^\\circ$, $170^\\circ<l<215^\\circ$, $|b|<5^\\circ$), 5 times more sensitive, in which we expect to detect thousands of galaxies (based on the HIMF of Davies et al., 2011) but for which observations are not yet complete.  This paper presents the discovery and the results from follow-up observations of a nearby galaxy, ALFA ZOA J1952+1428. Section 2 describes the discovery and follow-up with the Arecibo Radio Telescope. Section 3 describes follow-up observations with the Expanded Very Large Array\\footnote{The National Radio Astronomy Observatory (NRAO) is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (EVLA). Section 4 describes ongoing optical follow-up with the 0.9-m Southeastern Association for Research in Astronomy\\footnote{The SARA-North telescope at Kitt Peak National Observatory is owned and operated by the SARA consortium.} (SARA) telescope. Section 5 discusses the results from these observations. ",
        "conclusions": "\\subsection{Distance} ALFA ZOA J1952+1428 has a heliocentric velocity of $v_{hel}$ = +279 $\\mathrm{km \\; s^{-1}}$. Solving for its Local Group centered velocity using derivations of the solar motion with respect to the Local Group by Courteau $\\&$ van den Bergh (1999) gives $v_{LG}$ = 491 $\\mathrm{km \\; s^{-1}}$. Using Hubble's Law with $H_{0}$ = 70 $\\; \\mathrm{km \\; s^{-1} \\; Mpc^{-1}}$ puts this source at a distance of 7 Mpc. However, Hubble's Law is not a reliable distance indicator here because the dispersion of peculiar velocities in the local universe ($v_{hel} < 10,000$ $\\mathrm{km \\; s^{-1}}$) is $\\sigma = 298 \\pm 34$ km s$^{-1}$ (Masters 2008). The galaxy is probably not closer than 3 Mpc as the H\\,{\\sc{i}} linear size at this distance would be smaller than most compact galaxies containing H\\,{\\sc{i}} (Huchtmeier et al. 2007). For the following analysis, we take the distance to be 7 Mpc, although future observations may well revise this number.  \\subsection{Classification} As can be seen in the EVLA and SARA images, the H\\,{\\sc{i}} peak is slightly offset ($\\Delta\\theta = 8.7^{\\prime\\prime}$) from the optical emission, indicating either a false counterpart or a disturbed H\\,{\\sc{i}} distribution. The offset is $\\sim$300 pc at 7 Mpc, which is not uncommon even for isolated galaxies (c.f. $\\sim$400 pc offset in VV 124, Bellazzini et al. 2011). This could conceivably be a pair of low-surface brightness dwarf galaxies (c.f. HIZSS 3 with separation of $\\sim$900 pc, Begum et al. 2005), but there is no evidence for a second peak in the high signal-to-noise H\\,{\\sc{i}} spectrum shown in Figure 1. Further, ALFA ZOA J1952+1428 has half the velocity width that the pair in HIZSS 3 appeared to have; W$_{50}$ = 55 km s$^{-1}$ for HIZSS 3 (Henning et al. 2000) compared to W$_{50}$ = 28 km s$^{-1}$ here. Any second galaxy would have to be much closer both spatially and in velocity than the pair in HIZSS 3 in order to escape detection. Deeper interferometric observations would be needed to be entirely conclusive.  It is possible that ALFA ZOA J1952+1428 is a high velocity cloud (HVC) co-incident with an optical source (c.f. HIPASS J1328-30, Grossi et al. 2007), though this is unlikely as there is strong evidence that ALFA ZOA J1952+1428 is not an HVC. Its recessional velocity does not lie near HVCs in this part of the sky (c.f. Figure 3a in Morras et al. 2000). The nearest population of HVCs is the Smith Cloud which lies $10^{\\circ}$ and 170 $\\mathrm{km \\; s^{-1}}$ away (Lockman et al. 2008) at its nearest point. If ALFA ZOA J1952+1428 were an HVC, it would be a remarkable outlier.  Furthermore, the velocity field of ALFA ZOA J1952+1428 shows a gradient ten times larger than those of HVCs (Begum et al. 2010).  ALFA ZOA J1952+1428 appears to be a dwarf galaxy judging from its Gaussian H\\,{\\sc{i}} profile and low H\\,{\\sc{i}} mass. At a distance of 7 Mpc, $M_{HI}$ = $10^{7.0}$ M$_\\odot$, which is significantly lower than the gaseous content of spiral-type galaxies (Roberts and Haynes 1994). Also, its low luminosity ($L_{B}$ = $10^{7.5}$ L$_\\odot$ at 7 Mpc), H\\,{\\sc{i}} content, and blue colors are strong evidence that it is not an early-type galaxy.  There is no possible counterpart visible in 2MASS archive images or listed within 8$^\\prime$ in the 2MASS Extended Source Catalog (Jarrett et al. 2000). We plan follow-up NIR observations later this year with the 1.4-m Infrared Survey Facility in Sutherland, South Africa. We will use observations by its main instrument, SIRIUS, which has three detectors that operate simultaneously ($J$, $H$, $Ks$) with a field of view of 7.8 $\\times$ 7.8 arcmin.  Table 1 summarizes the observational data and derived quantities. Columns 1 and 2 give equatorial coordinates (J2000) for the  H\\,{\\sc{i}} peak. Columns 3 and 4 give the Galactic coordinates. Column 5 gives the heliocentric velocity from the mid-point of the velocity width at 50\\% peak flux. Column 6 gives the velocity width at 50\\% peak flux. Column 7 gives the integrated flux. Column 8 gives the M$_{HI}$/L$_B$ ratio using the m$_B$ calculated from the SARA telescope. The error on L$_B$ is dominated by the unknown uncertainty in A$_B$, thus we do not quote an error on M$_{HI}$/L$_B$. Column 9 gives the angular size of the  H\\,{\\sc{i}} out to the 1 M$_\\odot$ pc$^{-2}$ level. The last two Columns give values as a function of the distance to the galaxy in Mpc. Column 10 gives the linear size of the H\\,{\\sc{i}} at its largest extent. Column 11 gives total H\\,{\\sc{i}} mass.  \\begin{table*} \\caption[]{Properties of ALFA ZOA J1952+1428.} \\label{tab:mba} \\scriptsize \\begin{tabular}{ccccccccccc} R.A. & Decl. & $l$ & $b$ & $v_{hel}$ & $W_{50}$ &$F_{HI}$& M$_{HI}$/L$_B$ & H\\,{\\sc{i}} Size & $A_{o}$ & M$_{HI}$\\\\ (J2000.0) & (J2000.0) &  (deg) & (deg) & ($\\mathrm{km \\; s^{-1}}$) & ($\\mathrm{km \\; s^{-1}}$) &(Jy $\\mathrm{km \\; s^{-1}}$)& (M$_\\odot$/L$_\\odot$) & (arcsec) & (kpc) &(M$\\odot$)\\\\ \\tableline 19 52 11.8 & +14 28 24 &  52.82 & -6.42 & 279 $\\pm$ 1  & 28 $\\pm$ 2&0.94 $\\pm$ 0.07& 0.3 & 44 $\\times$ 40 & 0.21D &$10^{5.3}$D$^{2}$\\\\ \\end{tabular} \\end{table*}  \\subsection{Morphological Type} Compared to other dwarf galaxies (Roberts \\& Haynes 1994; O'Neil et al. 2000), ALFA ZOA J1952+1428 is not particularly gas rich, $M_{HI}/L_B$ = 0.3 M$_\\odot$/L$_\\odot$, but it is very small and blue with an H\\,{\\sc{i}} linear size of 1.4 $\\times$ 1.3 kpc at 7 Mpc and $B$ - $R$ = 0.1 mag. The H\\,{\\sc{i}} mass, $\\mathrm{M_{HI}/L_B}$ ratio, blue optical colors, and linear size of ALFA ZOA J1952+1428 are similar to those of blue compact dwarf (BCD) galaxies (Huchtmeier et al. 2007). BCDs are small, blue, irregular dwarf galaxies which have low surface brightness features, ongoing star formation, and higher metallicities than typical dwarf galaxies. The velocity field of ALFA ZOA J1952+1428 shows structure, but non-uniform rotation which is common in blue compact dwarf galaxies (Ramya et al. 2011). The velocity dispersion map shown in Figure 2 shows a significant amount of dispersion around the stellar looking object on the left side of the galaxy. This could be an ionized hydrogen region. H-alpha observations with the SARA telescope will examine this and quantify the star formation in this system. Deep $B$, $V$, and $R$ band observations with the SARA telescope will reveal whether there are low-surface brightness features in ALFA ZOA J1952+1428.  Alternatively, there is evidence for the existence of blue, metal-poor, gas-rich dwarf galaxies on the margins of galaxy groups (Grossi et al. 2007). These dwarfs are old (2 - 10 Gyrs) but have had remarkably little star formation in their history. They are thought to be galaxies in transition between dwarf irregular and dwarf spheroidal galaxies. ALFA ZOA J1952+1428 differs from the Grossi et al. galaxies because it has a lower $\\mathrm{M_{HI}/L_B}$ ratio and appears to be a field galaxy, though it may be a part of a group behind the Milky Way that has not yet been discovered.  There is a recently discovered Local Group galaxy, VV124 (Bellazzini et al. 2011), which is similar to ALFA ZOA J1952+1428 in size, H\\,{\\sc{i}} mass, and $M_{HI}/L_B$ ratio. This galaxy is isolated, as ALFA ZOA J1952+1428 appears to be. There is only one galaxy (3.8$^\\circ$ away) with $v_{hel} < 1000$ $\\mathrm{km \\; s^{-1}}$ within 10$^\\circ$ of ALFA ZOA J1952+1428 in NED (though this is not unusual in the ZOA). VV124 also shows an offset between the H\\,{\\sc{i}} and the optical counterpart as well as a velocity field with structure but non-uniform rotation. VV124 is considered to be a precursor of modern dwarf spheroidal galaxies that did not undergo an interaction-driven evolutionary path. Follow-up observations will reveal whether ALFA ZOA J1952+1428 has metallicity and star formation rates similar to a VV124-type galaxy.  \\subsection{Group Membership} It is possible that ALFA ZOA J1952+1428 is a Local Group galaxy, but it appears unlikely; ALFA ZOA J1952+1428 does not follow the relationship between radial velocity and angle from the solar apex that most other Local Group members do, as can be seen in Figure 3 (Courteau $\\&$ van den Bergh 1999). Further, the linear size of the H\\,{\\sc{i}} would be 210 pc at 1 Mpc, which is 4 times smaller than the smallest compact dwarfs (Huchtmeier et al. 2007).  \\begin{figure} \\plotone{fig3.eps} \\caption{Heliocentric velocity versus the cosine of the angular distance from the solar apex (modified with permission from Courteau and van den Bergh 1999). The black dots are Local Group galaxies and the red, open circle is ALFA ZOA J1952+1428. The solid line represents the solar motion solution of Courteau and van den Bergh ($v$ = 306 $\\mathrm{km \\; s^{-1}}$, $l = 99^\\circ$, $b = -3^\\circ$). The dotted lines represent the Local Group radial velocity dispersion, $\\sigma_{r} = $ 61 $\\mathrm{km \\; s^{-1}}$.} \\end{figure}  There are no known galaxy groups within $60^{\\circ}$ with $v_{hel} <$ 1000 $\\mathrm{km \\; s^{-1}}$ (Fouqu\\'{e} et al. 1992), making ALFA ZOA J1952+1428 either a field galaxy or a member of an undiscovered nearby group. Continued analysis of the ALFA ZOA data may clarify this."
    },
    "1107/1107.1085_arXiv.txt": {
        "abstract": "{In about half of Seyfert galaxies, the X-ray emission is absorbed by an optically thin, ionized medium, the so-called ``Warm Absorber'', whose origin and location is still a matter of debate. } {The aims of this paper is to put more constraints on the warm absorber by studying its variability.} {We analyzed the X-ray spectra of a Seyfert 1 galaxy, Mrk~704, which was observed twice, three years apart, by XMM-$Newton$.} {The spectra were well fitted with a two zones absorber, possibly covering only partially the source. The parameters of the absorbing matter - column density, ionization state, covering factor - changed significantly between the two observations.  Possible explanations for the more ionized absorber are a torus wind (the source is a polar scattering one) or, in the partial covering scenario, an accretion disk wind. The less ionized absorber may be composed of orbiting clouds in the surroundings of the nucleus, similarly to what already found in other sources, most notably NGC~1365. }{} ",
        "introduction": "Absorption from ionized matter in the X-ray spectrum of AGN (the so-called warm absorber) was discovered many years ago (Halpern 1984). Since then, many advances in its understanding have been made, and we know now that it is present in about half of Seyfert galaxies (e.g. Reynolds 1997) and that the matter is photoionized. Warm absorbers are also known to vary, and indeed the first discovered absorber was variable (Halpern 1984). The location of the warm absorber (or, indeed, of the absorbers, as more than one ionizing zone is often found) is, however, largely uncertain. There is some evidence for its origin as a wind from the dusty torus envisaged in unification models for Seyfert galaxies (Blustin et al. 2005), but cases in which an origin from the disk seems to be preferred do also exist (Krongold et al. 2007).  Mrk~704 is a local ($z$=0.029234) Seyfert 1.2 galaxy (Veron-Cetty \\& Veron 2010), bright enough in X-rays to be detected by $Swift$/BAT (Ajello et al. 2008). In this paper we report on extreme warm absorber variability, on yearly time scales, revealed by two XMM-$Newton$ observations, and possible variability on monthly time scales from short Swift/XRT observations.  The paper is organized as follows: in Section~2 we report on the XMM-$Newton$ observations and data reduction, while the relative data analysis are discussed in Section~3. Section~4 presents the analysis of the $Swift$ and $ASCA$ observations, while the results are summarized and discussed in Section~5.  \\begin{figure*} \\epsfig{file=lc_obs1.ps,width=6cm,angle=-90} \\hfill \\epsfig{file=lc_obs2.ps,width=6cm,angle=-90} \\caption{The 0.3-2 keV (upper panel), 2-10 keV (middle panel) and the hard-to-soft ratio (lower panel) light curves, for the first (left panels) and second (right panels) XMM-$Newton$ observations.} \\label{lc_obs} \\end{figure*} ",
        "conclusions": "The Seyfert 1.2 galaxy Mrk~704 was observed twice by XMM-Newton, about three years apart (first observation in October 2005, the second in November 2008).  While the hard ($>$5 keV) spectrum was almost constant between the two observations (with a possible broadening/strenghtening of the iron line in the first observation), the soft part dramatically changed, showing much heavier absorption in the first observation. A strong Neon IX line is also clearly present in the first observation, while it is not detected in the second one. The very similar hard X-ray spectrum in the two observations (the behaviour of the intrinsic soft X-ray spectrum is more uncertain because of the strong absorption), together with the change in the column density, suggest that the variations are due to a change in the properties of the absorbing clouds. Dramatic changes of the column density and/or the ionization state of the absorber have already been observed in other sources, the most spectacular being NGC 4151 (e.g. Schurch \\& Warwick 2002, de Rosa et al. 2007), NGC 7582 (Bianchi et al. 2009) and NGC 1365 (Risaliti et al. 2005, 2009). In the last two sources, in particular, large variations of the column density of the (cold) absorber occur on time scales as short as less than a day. Here we observe also a strong variation of the ionization state of the absorbers. As the primary continuum (at least the hard one) stayed almost constant, this indicates a variation of the location and/or of the density of the clouds.  It is interesting to note that Mrk~704 is a ``polar-scattering'' Seyfert 1, i.e. a source with optical polarization aligned perpendicularly to the radio source axis, as usually found in Seyfert 2s (Smith et al., 2004). Smith et al. suggest that in these sources the nucleus is seen through the edge of the torus.  We fitted the spectra with two absorbing regions, either fully or partially covering the primary emission. The improvement in the quality of the fit when letting  the aborbers be partial is only moderate for the first observation, but more significant in the second one. As the results are somewhat different in the two cases, let us discuss them separately.  With the full covering absorbers, both absorbing zones are found more ionized and with a lower column density in the second observation (even if the difference in the column density for the low ionization zone is only marginal). The RGS analysis of the first observation suggests a possible (but only marginally significant) inflow for the low ionization, low column density absorber, and an outflow velocity (very poorly determined, with 90\\% confidence level range between 1800 and 12000 km/s) for the high ionization, high column density absorber. An ionized and unstable torus wind, as suggested by Smith et al., may indeed provide an explanation for the latter zone. No significant inflow/outflow is instead found, for both absorbers, in the second observation.  The partial covering scenario, even if providing better fits, is more demanding geometrically, requiring a size of the obscuring clouds of the same order of the size of the emitting region. This naturally points to a much closer location of the obscuring clouds, as e.g. due to a radiatively-driven accretion disk wind (e.g. Proga 2003). Interestingly, the variability behaviour of Mrk~704 closely resemble that of mini-BAL QSOs (Giustini et al. 2010), strengthening the disk wind scenario, at least for the more ionized absorber (the colder absorber may be composed of orbiting clouds as those found in NGC~1365, Risaliti et al. 2009). Interestingly, a possible outflow has been detected in the first observation (but not in the second, despite very similar column densities and ionization parameters).  The alternative hypothesis, that the partial covering is mimicking the presence of a scattering component originating outside the absorbing region, is ruled out by the short term variability of the X-ray emission.  Mrk~704 was also observed four times by Swift/XRT between January 2006 and January 2007. Spectral and flux variability is apparent between the observations, but the quality of the spectra is not sufficient to establish the nature of the variations.  The iron K$\\alpha$ line, studied with XMM, is broad, much broader than the optical broad lines but narrower than expected if the entire line would be originated in the innermost accretion disk. However, if a narrow component (which seems to be almost ubiquitous in Seyfert galaxies) is added, then the remaining broad component is consistent with emission down to the innermost stable orbit, even if the quality of the data in not good enough to constrain the spin of the black hole.   \\begin{figure} \\epsfig{file=timeres.ps,width=6cm,angle=-90} \\caption{The spectra from the first and second part of the second XMM-$Newton$ observation, shown together for comparison.} \\label{timeres} \\end{figure}   \\begin{figure} \\epsfig{file=all_multi.ps,width=6cm,angle=-90} \\caption{The spectra of all XMM-$Newton$ and Swift observations shown together for comparison.} \\label{swiftspectra} \\end{figure}   \\begin{figure} \\epsfig{file=ascabestfit.ps,width=6cm,angle=-90} \\caption{Best fit model and data/model ratio for the $ASCA$ spectrum. See text for detail.} \\label{ascabestfit} \\end{figure}"
    },
    "1107/1107.0715.txt": {
        "abstract": " ",
        "introduction": "%%%%%%%%%%%%%%%%  Significant on-going efforts are being made to directly search for  Weakly Interacting Massive Particles (WIMPs).  An interesting and widely accepted signature of WIMPs is the annual modulation of their interaction rate, arising from the relative motion of the Earth around the Sun. Strikingly,  both the DAMA~\\cite{DAMA} and \\cogent~\\cite{cogent,Aalseth:2011wp} collaborations observe anomalous modulating events that may be interpreted as arising from interactions of spin-independent Dark Matter (DM).  As we discuss below, the two measurements are, to some extent,  consistent with each other, and point to a surprisingly low DM mass, of order a few GeV.  In contrast to the positive signals of DAMA and \\cogent, several other experiments find no evidence for DM.  Most notably, the CDMS~\\cite{CDMS}, {\\sc Xenon}10~\\cite{Angle:2009xb} and the \\xenon~\\cite{Xenon100} collaborations seem to disfavor  the parameter space indicated by DAMA and \\cogent. It is natural to ask, therefore, what possible systematic effects and/or DM properties can resolve the tension between the positive and null results?  One noticeable difference between DAMA/\\cogent{} and the null experiments, is that the latter veto electronic interactions while attempting to collect only  nuclear recoil events.  It is conceivable that the anomalous signals arise from such electronic recoils, a possibility that would explain away the existing tension.   A model of this type was considered in~\\cite{Kopp:2009et,Feldstein:2010su} prior to the recent \\cogent\\ measurement~\\cite{Aalseth:2011wp} and it remains to be seen whether this possibility is theoretically feasible. In this paper we pursue a different direction and study the viability of nuclear recoils of spin independent DM, as an explanation to the positive signals.  Several handles can, in principle, ameliorate the tension with the null results.  From the DM perspective one may consider, \\begin{itemize} \\item Inelastic scattering (endothermic or exothermic)~\\cite{TuckerSmith:2001hy,Graham:2010ca,Cline:2010kv}. \\item Isospin-violating couplings~\\cite{Kurylov:2003ra, Giuliani:2005my,Cotta:2009zu,Chang:2010yk,Kang:2010mh,Feng:2011vu,DelNobile:2011je}. \\item Velocity suppressed interactions. \\item Momentum dependent scattering~\\cite{Feldstein:2009tr,Chang:2009yt}. \\item Resonant scattering~\\cite{Bai:2009cd}. \\end{itemize} In addition there are uncertainties that may significantly change the expected scattering rates in various experiments: \\begin{itemize} \\item Astrophysical uncertainties~\\cite{Smith:2006ym,Bruch:2008rx,Fairbairn:2008gz,Kuhlen:2009vh,Fox:2010bz,Fox:2010bu,Savage:2006qr,Lisanti:2010qx,Arina:2011si}. \\item Possible  channeling effects~\\cite{Drobyshevski:2007zj,Bernabei:2007hw} (see however~\\cite{channeling}). \\item Uncertainties in the Xenon scintillation function at low recoil energies or in the statistical treatment of the background events (for a discussion, see e.g.~\\cite{Collar:2011wq}). \\end{itemize} % % A fully systematic analysis that takes into account all of the theoretical and experimental uncertainties above is hard to attain and we do not attempt here.   Instead,  we separately study the effects  of  most  of the above possibilities on the DM interpretation of the DAMA and \\cogent\\ results, as well as the null experiments.  Our goal is two-fold.  Primarily, we aim at understanding to what extent the DAMA and \\cogent\\ results are consistent with each other, and with the other null experiments (for related works, see~\\cite{Hooper:2011hd}).   In addition, we study what is required from the theoretical point of view and what needs to be assumed on the experimental side, in order to ameliorate the tension.  Along the way, we reanalyze the \\cogent\\ modulation data.  The paper is organized as follows.  In Section~\\ref{th} we summarize the standard computation of the expected signals and shortly discuss the DM velocity distributions considered here. In Section~\\ref{duccio} we describe the experimental data we use, dwelling on the various uncertainties and discussing their influence on the fits. Our results are presented in Section~\\ref{res}.  Here we follow the analyses endorsed by the various experimental collaborations and  explore whether velocity distributions or any of the DM scenarios mentioned above can explain the signals compatibly with the bounds. We conclude in section~\\ref{concl}.    %%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We explored the prospects of dark matter for jointly explaining the modulations signals observed  by DAMA and {\\cogent}  while complying with the constraints from other direct detection experiments: {\\sc Xenon10}, \\xenon, CDMS-Ge, CDMS-Si and SIMPLE. In the minimal scenario of spin-independent isospin-symmetric scattering there is a clear incompatibility, if we adhere to the data analyses performed by the experimental collaborations.  Taking into account uncertainties on the astrophysical velocity distribution of DM and  allowing for a higher sodium quenching factor in DAMA, $q_{\\rm Na}=0.4$ does not significantly improve the fits.  Furthermore, DAMA and {\\cogent} are mutually incompatible.  Assuming different DM cross sections on proton and neutrons allows to significantly improve the situation, making DAMA and {\\cogent} compatible and weakening one of the bounds from the null experiments. The best fit is for $f_p/f_n \\approx -1.5$, which allows to weaken the dominant bounds from  experiments with Xenon. However the overall global fit remains  poor, see Fig.\\fig{standardf}. Some additional improvement (although not very significant) is obtained by allowing  for a DM/nucleon form factor that depends on the transferred momentum and/or on the relative velocity. A similar results holds for inelastic scattering of DM with a mass splitting, $\\delta =M'_{\\rm DM}-M_{\\rm DM}$.  Both up-scattering ($\\delta > 0$) and down-scattering ($\\delta<0$) were considered.  The best global fit to all relevant experiments is obtained as in Fig.\\fig{iDM} for $|\\delta| \\approx 10$ keV.  A modest channeling in DAMA  allows a further slight improvement. Allowing also for a $\\sim 10\\%$ channeling in {\\cogent} results in a fit at smaller $M_{\\rm DM}\\approx 3\\GeV$, that, together with a $\\sim20\\%$ channeling in DAMA and inelasticity, allows a good global fit (Fig.\\fig{CH2}b).  We stress the alternative simple possibility of mild channeling in {\\cogent} alone: it allows for a good \\cogent\\ fit (without assuming a tuned inelasticity nor isospin-violating interactions) but excludes the DAMA signal.  \\medskip  In conclusion, we find it hard to explain the full set of current direct detection experiments with spin-independent scattering of dark matter, and contemplate on the possibility that one or both of the positive signals does not arise from DM interactions.  \\paragraph{Note added:}  While this work was being finalized, related works appeared~\\cite{Belli:2011kw,Schwetz:2011xm}.  \\paragraph{Acknowledgements} We thank Juan Collar, Mariangela Lisanti, Jeremy Mardon and Tracy Slatyer  for useful discussions. We wish to thank the authors of \\cite{weinerfit} for pointing out an error in Section 3.2 in an earlier version of the paper. The work of AS  was supported by the ESF grant MTT8 and by SF0690030s09 project.  The work of DP was supported by the Swiss National Science Foundation under contract No. 200021-116372. The work of MF was supported in part by the European Programme ``Unification in the LHC Era\",  contract PITN-GA-2009-237920 (UNILHC).  The work of TV was supported in part by the Director, Office of Science, Office of High Energy and Nuclear Physics, of the US Department of Energy under Contract DE-AC02- 05CH11231. AS thanks the invitation from the Berkeley physics department, where this work was initiated. This work was supported by the EU ITN \u00d2Unification in the LHC Era\u00d3, contract PITN- GA-2009-237920 (UNILHC).   \\footnotesize"
    },
    "1107/1107.4519_arXiv.txt": {
        "abstract": "We present the space spectrometer \\pam\\ observations of proton and helium fluxes during the December 13 and 14, 2006 solar particle events. This is the first direct measurement of the solar energetic particles in space with a single instrument in the energy range from $\\sim$ 80 MeV/n up to $\\sim$ 3 GeV/n. In the event of December 13 measured energy spectra of solar protons and helium were compared with results obtained by neutron monitors and other detectors. Our measurements show a spectral behaviour different from those derived from the neutron monitor network. No satisfactory analytical fitting was found for the energy spectra. During the first hours of the December 13 event solar energetic particles spectra were close to the exponential form demonstrating rather significant temporal evolution. Solar He with energy up to ~1 GeV/n was recorded on December 13.  In the event of December 14 energy of solar protons reached ~600 MeV whereas maximum energy of He was below 100 MeV/n. The spectra were slightly bended in the lower energy range and preserved their form during the second event. Difference in the particle flux appearance and temporal evolution in these two events may argue for a special conditions leading to acceleration of solar particles up to relativistic energies. ",
        "introduction": "The \\pam\\ spectrometer \\citep{picozza_pamela_2007} is a space instrument designed for the study of primary charged particles and antiparticles in a wide energy interval, mainly from tens of MeV to ${\\approx}$ 1.2 TeV for protons. It was launched in an elliptical orbit at an altitude between 350 and 610 km and inclination of 70 degree in June 2006. The main scientific goal of \\pam\\ is the measurement of the particle and antiparticle component in galactic cosmic rays \\citep{adriani_anomalous_2009, adriani_new_2009}, the study of galactic cosmic ray modulation by solar activity, and solar energetic particles \\citep{casolino_cosmic-ray_2006,de_simone_study_2009}. This paper reports the \\pam\\ measurements of the solar proton and helium fluxes in the energy range from below 100 MeV/n to several GeV/n during the December 13 and 14, 2006 solar particle events.   The problem concerning mechanism and site of solar energetic particles (SEP) acceleration remains an open question. Certainly, SEP may be produced after powerful explosive events on the Sun, accompanied by solar flares, Coronal Mass Ejections (CME), bursts of solar X /gamma-rays and radio emission \\citep{reames_particle_1999}. It is clear that not a single mechanism is involved into the SEP generation. Stochastic acceleration, shock acceleration, and acceleration by the DC electric fields in the process of magnetic reconnection are the main candidates. Acceleration of SEP may take place in the flare region, solar corona and even in the interplanetary space. It should be kept in mind that SEP themselves, while propagating, create conditions for the energy redistribution \\citep{tylka_new_2001,lee_coupled_2005}. Therefore, the energy spectrum of SEP provides valuable information for study of solar and interplanetary plasma processes. Most energetic particles are usually (not always) accelerated during short time close to the energy explosion, thus the first particles arriving at the observer site keep more information about the primary acceleration process. Less energetic SEP may be accelerated during hours and even days. Conditions for particle acceleration in the corona and interplanetary space are determined by temporal and spatial evolution of a shock \\citep[e.g.][]{ellison_shock_1985}.  Since the energy range of SEP extends over more than 5 orders of magnitude, and the SEP fluxes occupy more than 8 orders of magnitude, the SEP energy spectrum has never been measured by a single instrument. The bulk of the SEP observations has been performed on spacecraft and covers the particle energy below several hundreds MeV/n, whereas the relativistic solar particles have been observed by ground-based installations and remain less elucidated. This kind of observations and associated modeling can also be used to issue radiation alerts to be used by manned and unmanned missions. A multi-detector study of the solar particle event occurring in 2005 and the corresponding  GLE69 was performed on-ground using NM and SREM  radiation environment units in space \\cite{astra-7-1-2011}  It is not yet clear if relativistic solar particles require a special scenario for their generation. Some authors believe that relativistic solar protons have two components - the fast and the delayed ones. The fast component is initially accelerated in the flare region \\citep[e.g.][]{vashenyuk_relativistic_2006,vashenyuk_characteristics_2008,vashenyuk_relativistic_2008,mccracken_two_2008}, probably via magnetic field reconnection, whereas another mechanism, probably shock or stochastic acceleration, generates the second particle component. \\cite{bieber_spaceship_2004,bieber_relativistic_2005} argue against two components and in favour of the CME-driven shock as a main accelerator of the whole SEP population. A wide variety in concomitant acceleration conditions supports the point that various mechanisms play the main role in different SEP events \\citep{bombardieri_improved_2008,bazilevskaya_early_2009}. Detailed studies of more solar particle spectra in wide energy range are needed to resolve this problem.    Since ground-based instruments can only detect secondary cosmic rays, knowledge of response to the primary particles is necessary in order to find SEP fluxes in space. Some methods were developed \\citep{shea_possible_1982,cramp_october_1997,bieber_spaceship_2004,plainaki_modeling_2007,vashenyuk_relativistic_2006} to derive energy spectra and angular anisotropy of solar particles using the enhancements in the count rates of neutron monitors. Nevertheless, the particle fluxes derived from the neutron monitor data are model dependent.  Before the \\pam\\ launch, direct measurements of the relativistic solar particles were not fulfilled. The aim of this paper is to present the results of the first direct measurements by a single instrument of the solar protons in the energy range from 80 MeV to several GeV and the helium nuclei from 75 MeV/n up to $\\sim$ 700 MeV/n during the SEP events of December 13 and 14, 2006. The absolute intensities and energy spectra are compared with the results of direct (GOES, ACE) and indirect (neutron monitors, IceTop) measurements. ",
        "conclusions": "The \\pam\\ spectrometer was the first instrument which directly measured the relativistic SEP in the near Earth space. It is important since all previous such measurements were fulfilled with the ground-based installation and the derived SEP fluxes depended on the instrument response functions. The spectra of solar protons in the energy range 80 MeV - 3 GeV and helium 75 MeV/n - 1 GeV/n were measured during the first polar \\pam\\ passage at 0318-0329 UT. There is a good agreement between the protons fluxes measured by \\pam\\ and those obtained by the IceTop installation \\citep{abbasi_solar_2008,kuwabara_t._private_2009}. Keeping in mind accuracy of estimation of absolute SEP intensities from the neutron monitor data \\citep{vashenyuk_relativistic_2006,plainaki_modeling_2007,bombardieri_improved_2008,NMBANGLE2009AdSpR..43..474P}, reasonable agreement can be stated between the \\pam\\ and neutron monitor SEP fluxes. However, the \\pam\\ spectra are always harder in the low-energy interval probably indicating that neutron monitor yield functions are underestimated below $\\sim$700 MeV. During the second polar \\pam\\ passage the difference between the SEP fluxes taken from \\pam\\ and neutron monitors has become larger while agreement between \\pam\\ and IceTop remained very good. The correctness of the \\pam\\ observations was also confirmed by agreement with the direct SEP measurement on a balloon in the atmosphere.  Evolution of intensity and spectral shape of relativistic solar protons is often used to derive some information about SEP generation \\citep{bombardieri_improved_2008,vashenyuk_relativistic_2006,vashenyuk_characteristics_2008,vashenyuk_relativistic_2008,mccracken_two_2008, moraal_analysis_2008}. The indications were found that the first arriving relativistic particles would be accelerated in the flare region and have an exponential spectrum whereas the latter particles would be accelerated by a CME-driven shock and have a power-law spectrum. In this case an appropriate dynamics of the energy spectrum could be observed by \\pam. Being at the low latitudes \\pam\\ missed the earliest anisotropic phase of the event of December 13. The results of the first \\pam\\ observation around 0320 UT on December 13 confirmed existence of a hard quasi-exponential spectrum expected from the magnetic reconnection in the flare region. However, in spite of changes in the SEP intensity a quasi-exponential spectrum persisted, actually till the advent of the newly generated protons of the December 14 event. It should be noted that in the energy interval relevant to the neutron monitor observations ($>$ 500 MeV) the spectrum could be fitted by a power law beginning from $\\sim$1000 UT, however at lower energies the spectrum was flatter. We did not find any spectral form which would fit the observed spectra satisfactorily in the whole energy range and could prove a certain dominating mechanism of SEP acceleration. A quasi-exponential spectrum and fast temporal evolution of particle fluxes during several hours were present only in the event of December 13 when the relativistic protons were generated. In the event of December 14, without relativistic particles, a form of the solar proton energy spectra changed little and was almost power-law throughout the event. In addition, no Helium with energy above $\\sim$100 MeV/n was observed in the December 14 event. This may be indicative of special conditions leading to acceleration of particles up to relativistic energy."
    },
    "1107/1107.2109_arXiv.txt": {
        "abstract": "{We introduce  the Gauge Vector-Tensor (GVT) theory  by extending the AQUAL's approach  to  the GravitoElectroMagnetism (GEM) approximation of gravity.   GVT is a generally covariant theory of gravity  composed of a pseudo Riemannian metric and two $U(1)$ gauge connections that reproduces  MOND in  the limit of very weak gravitational fields  while remains consistent with the Einstein-Hilbert gravity in the limit of strong and Newtonian gravitational fields.   GVT also provides a simple framework to study the GEM approximation to gravity. We illustrate that  the  gravitomagnetic force at the edge of a galaxy can  be in accord with either GVT  or $\\Lambda$CDM but not both.   We also study the physics of the GVT theory around  the  gravitational saddle point of the Sun and Jupiter system. We notice that the conclusive  refusal of  the GVT theory demands measuring    either  both of the gravitoelectric and gravitomagnetic fields inside the Sun-Jupiter  MOND window, or the gravitoelectric field inside two different solar GVT MOND windows. The GVT theory, however, will be favored by observing an anomaly in the gravitoelectric field inside a single MOND window. }   \\begin{document} ",
        "introduction": "Either $95\\%$ of the observed universe is made of things that have not yet been observed in the Solar system, or the law of dynamics or gravity should be modified in  very low accelerations or very weak gravitational fields. The  $\\Lambda$-CDM model of cosmology buys the first approach. Its challenges \\cite{Famaey:2011kh,Kroupa:2013yd}, however,  signal that \"the physics of the dark sector is, at the very least, much richer and complex than currently assumed, and that our understanding of gravity and dynamics might also be at play\" \\cite{Famaey:2013ty}.   The second approach is the modified theories of gravity. The modified theories of gravity can be classified into the following two categories: \\begin{enumerate} \\item Phenomenological search for  the dynamics of the metric. \\item Introducing new degrees of freedom for gravity in addition to the metric. \\end{enumerate} The first class assumes that gravity is described by a  pseudo Riemannian metric and the action of gravity is given by \\begin{equation} S = \\int \\sqrt{-\\det g} \\big({\\cal L}_g[g_{\\mu \\nu}, R_{\\mu\\nu\\lambda\\eta}, \\nabla_\\mu]+{\\cal L}_m[\\Psi]\\big)\\,, \\end{equation} where $R_{\\mu\\nu\\lambda\\eta}$ is the Riemann tensor constructed out of the metric $g_{\\mu\\nu}$, ${\\cal L}_m[\\Psi]$ is the matter's action not necessarily minimally coupled to the metric, and ${\\cal L}_g$ is the gravity's action. The Einstien-Hilbert theory assumes ${\\cal L}_g=R$ where $R$ is the Ricci scalar. The purchasers of   this class choose to reject the  Einstein-Hilbert assumption and search for families of ${\\cal L}_g$ reproducing the dynamics of nature in large scales. Considering the infinite number of possibilities in choosing $\\cal L$ and the finite set of the cosmological data, this purchase will work \\cite{Exirifard:2008dy}. It would  not necessarily be in accord  with the principle of the Occam's razor.  It also will lead to a set of nonlinear partial differential equations of degrees larger than two, a set of equations which most of  rational humans  would  despise. These are, however, the prices to pay.  The second class of the modified theories of gravity introduces new degrees of freedom in addition to the metric to describe gravity. The most known example of this class is the TeVeS theory  \\cite{Bekenstein:2004ne}. TeVeS introduces a pseudo Riemannian metric, a scalar and a vector field in order to phenomenologically describe the physics  in very weak gravitational fields (the MOND regime) . The TeVeS theory defines new nonlinearity in order to solve the physics of the MOND regime. The introduced  nonlinearity, however, is not local to the MOND regime of the theory. It continues to the very strong gravity regime of theory.   The physics of very strong gravitational systems, therefore,  strongly constrain the TeVeS theory. This signals that the introduced nonlinearity of the TeVeS is not appropriate to describe the physics of the MOND regime. One should define a nonlinearity capable of producing the physics of the MOND system such that  the nonlinearity does not propagate all the way down to the Newtonian and strong regime of the theory. In order to  perform such a definition, we go back to  the very root of the TeVeS theory: the AQUAL theory. We show how to apply the AQUAL procedure upon the GravitoElectroMagnetism approximation of gravity. This introduces a non-covariant version for GEM in MOND regime whose generally covariant version demands introducing gauge vector fields rather than a scalar field. We, thus, introduce two $U(1)$ gauge vectors in addition to the metric and present a generally covariant theory for GEM in the MOND regime. This theory, which we call the Gauge Vector Tensor theory, reproduces  MOND in  the limit of very weak gravitational fields  while remains consistent with the Einstein-Hilbert gravity in the limit of strong and Newtonian gravitational fields.   In contrary to the TeVeS theory,  the GVT theory is in total agreement with the physics of the strong gravity. Its equations of motion are also much simpler than those of the TeVeS theory.  The paper is organized as follows: Sections \\ref{subsectionGEM1} and \\ref{subsectionGEM2}  review the GravitoElectroMagnetism (GEM) to gravity.   Section \\ref{reviewAQUALmodel} reviews the algorithm that leads to the AQUAL theory as a realization of the MOND paradigm. Section \\ref{upliftAQUAL}  applies the AQUAL's algorithm to  GEM.  Section \\ref{section6} introduces one gauge field and presents a covariant realization of the GEM to MOND. It also discusses the phenomenological constraints on the theory. Section \\ref{section7} introduces an additional gauge field in order to make the theory fully consistent with the predictions of the Einstein-Hilbert theory for the strong and Newtonian gravitational field. Section \\ref{section8} studies various regimes of the GVT theory.  The GVT theory possesses the Newtonian and strong regime of gravity, the MOND regime and the post-MONDian regime. Section \\ref{section9} calculates the gravitomagnetic field of a spinning galaxy in the $\\Lambda$CDM theory and the GVT theory. It shows that the gravitomagnetic field of a galaxy can be in accord with only one of them.   Section \\ref{fasle4-2} studies the physics of the GVT theory around  the  gravitational saddle point of the Sun and Jupiter system. It notices that the conclusive  refusal of  the GVT theory demands measuring    either  both of the gravitoelectric and gravitomagnetic fields inside the Sun-Jupiter  MOND window, or the gravitoelectric field inside two different solar GVT MOND windows. It concludes that the GVT theory, however, can be favored by observing an anomaly in the gravitoelectric field inside a single MOND window. Section \\ref{section11} provides the conclusion and outlook. ",
        "conclusions": "\\label{section11} We have introduced the Gauge Vector Tensor theory: a generally covariant theory of gravity   composed of a pseudo Riemannian metric and two $U(1)$ gauge connections that reproduces  MOND in  the limit of very weak gravitational fields  while remains consistent with the Einstein-Hilbert gravity in the limit of strong and Newtonian gravitational fields. The nonlinearity introduced by the GVT theory to reproduce the MOND behavior resides only inside the MOND regime and it does not propagates to the strong regime of gravity. This is a clear advantage of the GVT theory over the Bekenstein's Tensor-Vector-Scaler theory \\cite{Bekenstein:2004ne}. We have been motivated to introduce the GVT theory after uplifting  the GravitoElectroMagnetism approximation to gravity to the Milgrom's MOND theory \\cite{MOND}.  We have illustrated that  the  gravitomagnetic force at the edge of a galaxy can  be in accord with either GVT  or $\\Lambda$CDM but not both.   We also have studied the physics of the GVT theory around  the  gravitational saddle point of the Sun and Jupiter system. We have noticed that the conclusive  refusal of  the GVT theory demands measuring    either  both of the gravitoelectric and gravitomagnetic fields inside the Sun-Jupiter  MOND window, or the gravitoelectric field inside two different solar GVT MOND windows. The GVT theory, however, can be favored by observing an anomaly in the gravitoelectric field inside a single MOND window.  We also need to study the cosmology and the gravitational lensing of the GVT theory.   Let it be hasten that, as shown in section \\ref{section6},  the GVT theory is an extension of the Moffat's Scalar-Tensor-Vector theory \\cite{Moffat:2005si}. We, therefore, envisage that it inherits most of the merits of the Moffat's theory  in describing the gravitational lensing and cosmology.   We, however, accomplish this study elsewhere."
    },
    "1107/1107.5729_arXiv.txt": {
        "abstract": "We propose two new means of identifying the main class of progenitors of Type~Ia supernovae---single or double degenerate: (i)~If the range of supernova properties is significantly determined by the range of viewing angles of non-spherically symmetric explosions, then the nature of the correlation between polarization and another property (for example, the velocity gradient) can be used to determine the geometry of the asymmetry and hence the nature of the progenitor, and (ii)~in the double- but not in the single-degenerate case, the \\textit{range} in the observed properties (e.g., velocity gradients) is likely to increase with the amount of carbon seen in the ejecta. ",
        "introduction": "The fact that we are still uncertain about the nature of the progenitor systems of Type~Ia supernovae (SN\\,Ia)---some of the most powerful explosions in the universe---has become a major embarrassment for modern astrophysics. The problem is compounded by the fact that Type~Ia supernovae (SN\\,Ia) have become a key tool for placing constraints on the nature of the dark energy that propels the accelerating cosmic expansion (Riess et~al.\\ 1998; Perlmutter et~al.\\ 1999). The use of SN\\,Ia for the determination of such cosmological parameters requires the taming of evolutionary effects---an understanding, or at least a calibration, of the evolution of SN\\,Ia luminosity with cosmic time. A full understanding of these effects must hinge on a proper comprehension of the nature and properties of the progenitors (see, e.g., Livio 2000; Nomoto et~al.\\ 2000; Hillebrandt \\& Niemeyer 2000 for reviews, and Riess \\& Livio 2006 for a more recent discussion).  In the present very brief letter we present two simple \\textit{ideas} for observations and analyses that, while not easy, could be used to identify the dominant (in terms of its contribution to the SN\\,Ia rate) class of progenitors. ",
        "conclusions": "We have proposed two new ideas for observations that could potentially identify the main class of progenitors of SN\\,Ia---either single degenerate or double degenerate. Clearly, to allow for a detailed comparison between theory and observations, full-scale 3D simulations that include all the relevant processes will be required.  However, the results of such simulations will also necessarily depend on the input parameters and physics. Here we believe that we have identified robust trends, that can be used as distinguishing characteristics  between the two scenarios.  As excellent observations of velocity gradients are becoming available (e.g., Foley, Sanders \\& Kirshner 2011), coupling those with polarization measurements and measurements of the amount of carbon should become feasible in the not too distant future."
    },
    "1107/1107.0560_arXiv.txt": {
        "abstract": "Noether symmetry for Gauss-Bonnet-Dilatonic interaction exists for a constant dilatonic scalar potential and a linear functional dependence of the coupling parameter on the scalar field. The symmetry with the same form of the potential and coupling parameter exists all in the vacuum, radiation and matter dominated era. The late time acceleration is driven by the effective cosmological constant rather than the Gauss-Bonnet term, while the later compensates for the large value of the effective cosmological constant giving a plausible answer to the well-known coincidence problem. ",
        "introduction": "Many alternative theories of gravity have been proposed so far in order to present viable cosmological models of dark energy associated with observed cosmic acceleration. Among them, the introduction of a Gauss-Bonnet term into the Gravitational action has received much attention in recent years \\cite{a1}\\cite{a2}\\cite{a3}\\cite{a4}\\cite{a5}\\cite{a6}\\cite{a7}\\cite{a8} \\cite{a9}\\cite{a10}\\cite{a11}\\cite{a12}\\cite{a13}\\cite{a14}\\cite{a15}\\cite{od1},\\cite{od2}. In particular, important issues like - late time dominance of dark energy after a scaling matter era and thus alleviating the coincidence problem, crossing the phantom divide line and compatibility with the observed spectrum of cosmic background radiation have also been addressed recently \\cite{b1}\\cite{b2}. Gauss-Bonnet term arises naturally as the leading order of the $\\alpha ^{\\prime}$ expansion of heterotic superstring theory, where, $\\alpha^{\\prime}$ is the inverse string tension. \\cite{c1}\\cite{c2}\\cite{c3}\\cite{c4}.\\newline However, Gauss-Bonnet is a topologically invariant term in 4-d space-time and so it is coupled with a dilatonic scalar field\\textbf{,} in order to avoid collapse of the equations to those corresponding to standard cosmological model. As a result, at least two unknown functions are to be postulated, or derived, viz., the potential of the scalar field $V(\\phi)$ and the coupling of the Gauss-Bonnet term with gravity $\\Lambda(\\phi)$. A most elegant procedure is to make a single postulate in order to derive these two functions rather than setting both of them arbitrarily by hand. This may be done by demanding Noether symmetry amongst field variables, which has received much attention in recent times, particularly in the context of higher order theory of gravity \\cite{d1}\\cite{d2}\\cite{d3}.\\newline\\noindent The Noether symmetry approach for the solution of the cosmological equations was developed many years ago \\cite{e1}, and since then applied to find general exact solutions of many problems in the field \\cite{e2}\\cite{e3}\\cite{e4}\\cite{e5}\\cite{e6}\\cite{d1}. It consists first in recognizing that the field equations, when turn out to be ordinary differential equations, may be derived from an ordinary point Lagrangian. Then, it is required to select the (not yet established) functions under the condition that the Lagrangian should be preserved under some infinitesimal point transformation (Lie derivative). Once the functions are obtained, general exact integration of the field equations may be usually performed. If not, it simplifies the set of differential equations considerably, as in the present case, which helps in discussing the solutions and setting the values of parameters of the theory. The discussion on the physical implications of this symmetry may be found in  \\cite{e2}. As a matter of fact, its nature remains obscure, but it revealed so fruitful in may circumstances  that it is worth to attempt its applicability hereto. Of course, here like earlier works, the same results may be obtained by suitable guess of the functions and transformations. However it should be made clear that it is extremely difficult to make such a guess without Noether symmetry approach.   In the present work, our starting point is the gravitational action with Gauss-Bonnet term being coupled with a dilatonic scalar, in the presence of cold dark matter, for which Noether symmetry has been explored in the background of spatially flat Robertson-Walker metric. Noether symmetry has been found in the matter dominated era, after setting the state parameter $w$ corresponding to the baryonic and the cold dark matter to zero. In the process the potential has been fixed to a constant and the Gauss-Bonnet coupling parameter $\\Lambda(\\phi)$ turned out to be a linear function of $\\phi$. In the subsection 2.1, we have generated a set of solutions simply by handling the algebraic equation in Hubble parameter rather than solving differential field equations. The results thus obtained are intriguing.\\newline The original idea to include Gauss-Bonnet term into the action was to drive the late time cosmic acceleration, playing thus the key role of effective \"dark energy\", instead of the scalar field. On the contrary, we discovered that, at least in the circumstances described below, it is again the scalar field which drives the acceleration. The Gauss-Bonnet term instead plays the role of contrasting the effective cosmological constant, ie., nullifying the effective cosmological constant as should be clear below. \\ However, to reduce the cosmological constant by some $120$ order magnitude requires $\\Omega_{\\phi}$ of the same order of magnitude. Though there has been some early attempts in this regard \\cite{r1,r2}, nevertheless it is an interesting issue, since, \\ it has been observed that modified gravity theory can contrast the effective cosmological constant. In section 3 we compare our model with $\\Lambda$CDM, which as usual, shows that it is practically impossible to identify the two, as far as luminosity distance versus redshift graph is concerned. It is worth noting that the Noether symmetry approach could be applied also in the case of modified GB gravity theories, which includes, for instance, the functional dependence from the Gauss-Bonnet invariant G in the form of f(G) only, or also an additional dependence on curvature as f(G,R). Moreover, since Gauss-Bonnet is not a topologically invariant term in dimensions greater than 4, so we could also consider the standard Gauss-Bonnet gravity with or without dilatonic coupling (\\cite{od3}). But then, it should be clear, however, that the mathematical feasibility of the method will be different and of course very difficult. ",
        "conclusions": "Noether symmetry has been enforced in Gauss-Bonnet dilatonic scalar theory of gravity and the following important results have emerged.  \\noindent  \\begin{description} \\item[ (1) ] The coupling parameter and the potential thus found are independent of the evolutionary history of the Universe, ie., Noether symmetry exists throughout the history of evolution of the Universe starting from the early vacuum dominated era, passing over to radiation dominated era and finally at the matter dominated era. Existence of such symmetry fixes up the dilatonic-Gauss-Bonnet coupling parameter $\\Lambda=\\lambda\\phi$ and the scalar potential $V=V_{0}$ (a constant) once and forever.\\newline  \\item[ (2) ] The same form of $\\Lambda(\\phi)$ was obtained in an earlier work \\cite{b2} at the late time cosmic evolution, when the kinetic term viz., $\\dot\\phi^{2}$ indeed became constant and the Universe sets off for an accelerated expansion.\\newline  \\item[ (3) ] Since the potential is obliged to be a constant, the effective cosmological constant is now $\\Lambda_{eff}=\\Lambda_{eff(\\varphi)}% +\\Lambda_{GB}$, which may be comparable with the present Hubble parameter.\\newline  \\item[ (4) ] The late time cosmic acceleration is driven by the scalar field effective cosmological constant rather than the Gauss-Bonnet term.\\newline  \\item[ (5) ] Figure (2) depicts that the presence of the Gauss-Bonnet-dilatonic interaction term $\\lambda$ puts up the possibility to reduce the tremendous repulsive power of a large effective cosmological constant. Thus the  well known coincidence problem viz., why the cosmological constant is so small today might have been given a plausible answer.\\newline  \\item[ (6) ] For late universe, there is an almost perfect equivalence with the $\\Lambda$CDM model as depicted in Figure (4).  \\item[ (7) ] A more accurate fit procedure gives even more satisfactory results, with a best fit value for $\\Omega_{m}$, whose lower limit is consistent with the results obtained from very different kinds of methods using various astronomical objects, within the framework of the standard $\\Lambda CDM$ model \\cite{WMAP3,sab}. \\end{description}  {\\bf Acknowledgements:}  A.K. Sanyal is grateful to the University of Naples (Ufficio Relazioni Internazionali) for supporting a visit in the Department of Physical Sciences, Naples."
    },
    "1107/1107.1346_arXiv.txt": {
        "abstract": "Low luminosity gamma-ray bursts (\\llGRBs) constitute a sub-class of gamma-ray bursts (GRBs) that plays a central role in the GRB-supernova connection. While {\\llGRBs} differ from  typical long GRBs (LGRBs)  in many aspects, they also share some common features.  Therefore, the question whether the gamma-ray emission of {\\llGRBs} and LGRBs has a common origin is of great interest. Here we address this question by testing whether {\\llGRBs}, like LGRBs according to the Collapsar model, can be generated by relativistic jets that punch holes in the envelopes of their progenitor stars. The collapsar model predicts that the durations of most observed bursts will be comparable to, or longer than, the time it takes the jets to breakout of the star. We calculate the jet breakout times of {\\llGRBs} and compare them to the observed durations. We find that there is a significant access of {\\llGRBs} with durations that are much shorter than the jet breakout time and that these are inconsistent with the Collapsar model. We conclude that the processes that dominate the gamma-ray emission of {\\llGRBs} and of LGRBs are most likely fundamentally different. ",
        "introduction": "According to the Collapsar model \\citep{Paczynski98,MacFadyen99} the core collapse of a massive star results in the formation of a compact object, a black hole or a rapidly rotating neutron star. The compact object ejects a relativistic bipolar, baryon poor jet, along its rotation axis.  % The jet  punctures the surrounding stellar envelope and it emits the observed $\\gamma$-rays at a large distance from the star where the optical depth is small and the high energy photons can escape. This model, that  is accepted as the standard model for long GRBs (LGRBs),  explains naturally the association of some LGRBs with SNe, and their general emergence in star forming regions  \\citep[see][for recent reviews]{Woosley06, Hjorth11}.   A closer look at the members of the  spectroscopically confirmed GRB/SN group (on which the GRB/SNe association hinges) reveals that four out of the six detected bursts: GRB980425 (SN1998bw), GRB031203 (SN2003lw), GRB060218 (SN2006aj) and GRB100316D (SN2010bh), are quite different  than ``normal\" LGRBs (see \\S 2):  They are less luminous; have a smooth non-variable lightcurve, and show no evidence for a high energy power-law tail. Although only a handful of such bursts were observed, the small observable volume set by their low luminosity implies an event rate much higher than the rate of LGRBs  pointing towards Earth \\citep{Coward05,Cobb06,Pian06,Soderberg06,Liang07,Guetta07, Fan11}. The unique characteristics of these low luminosity GRBs (denoted hereafter {\\llGRBs}) suggest that they may be generated by a totally different process than most LGRBs. As such it is of great interest to check whether % {\\llGRBs} % can arise from Collapsars. Specifically  we ask the question: can {\\llGRBs} be generated by relativistic jets that break out of their progenitor stars.  To answer this question we study, in \\S 3, (following Bromberg et al. 2011; hereafter B11) the propagation of a relativistic jet in a stellar envelope.  We obtain the minimal conditions required for the jet to break out of the star. % Specifically, we estimate the minimal time that the central engine must power the jet for a successful crossing of the star. Using this minimal breakout time we examine the expected duration distributions of LGRBs, SGRBs and {\\llGRBs} (\\S 4). We  discuss the implications of this distribution on the origin of {\\llGRBs} in \\S 5 and we summarize our results  in \\S 6. ",
        "conclusions": "The activity of the central engine, within the Collapsar model, is independent of the breakout time of the jet. This implies that the duration of the majority of bursts should  be comparable to, or longer than, their jet breakout time and that for $T_{90}\\ll t_B$ the bursts' duration distribution should be constant, independent of $T_{90}$. These two predictions are satisfied by LGRBs, providing another support to the Collapsar model \\citep{B11b}. As expected, they are not satisfied by SGRBs that are not produced by Collapsars.  Like SGRBs, the observed distribution of {\\llGRBs}  is inconsistent with these  two predictions of the collapsar model. A  large fraction of {\\llGRBs} has $T_{90}/t_B\\ll 1$.  The probability that the observed {\\llGRBs} $T_{90}/t_B$ distribution is consistent with the LGRBs distribution is smaller than $5\\%$. Similarly, the {\\llGRBs} $T_{90}$ distribution  (for $T_{90}\\ll t_B$) is not flat at a confidence level  $>97\\%$.   Taken together with their peculiar $\\gamma$-ray emission, our overall conclusion % is that {\\llGRBs} are very unlikely to be produced by Collapsars like LGRBs. This can be reconciled with the fact that {\\llGRBs} are accompanied by ``engine driven\" SNe, if {\\llGRBs} are produced by ``failed jets\" that don't break out of their progenitors. These ``failed\" jets  deposit their energy into the stellar envelopes and the $\\gamma$-ray emission  arises from   shock breakout \\citep{Kulkarni98,MacFdyen01,Tan01,Campana06,Wang07,Waxman07,Katz10,Nakar11}. Our analysis doesn't prove this model but it strongly disfavors its major competitor, any variant on the Collapsar model.  The conclusion that the {\\llGRBs} originate from \"failed jets\" rather then successful ones has some interesting implications. In particular the high rate of {\\llGRBs} implies that jets which are generated in SNe have a higher chance of remaining buried than to break out. Accordingly most  SNe  engines can   generate jets that produce {\\llGRB} and only a few are powerful enough to produce jets that break out and produce  LGRBs \\citep{Mazzali08}.     We thank to Re'em Sari for helpful discussions. The research was supported by the Israel Center for Excellence for High Energy Astrophysics and by an ERC advanced research grant (OB and TP) and by by the Israel Science Foundation (grant No.\\ 174/08) and by an EU International Reintegration Grant (EN)."
    },
    "1107/1107.3175_arXiv.txt": {
        "abstract": "We use a numerical code  called PAKALMPI  to compute  synthetic spectra of the solar emission in quiet conditions at millimeter, sub-millimeter and infrared wavelengths. PAKALMPI solves the radiative transfer equation, with Non Local Thermodynamic Equilibrium (NLTE),  in a three dimensional % geometry using a multiprocessor environment. % The code is able to use three opacity functions: classical bremsstrahlung, $H^-$ and inverse bremsstrahlung. In this work we have computed and compared two synthetic spectra, one in the common way: using bremsstrahlung  opacity function and considering a fully ionized atmosphere; and a new one considering bremsstrahlung,  inverse bremsstrahlung and $H^-$ opacity functions in NLTE. We analyzed in detail the local behavior of the low atmospheric emission at 17, 212, and 405 GHz (frequencies used by the Nobeyama Radio Heliograph and the Solar Submillimeter Telescope). We found that the $H^-$ is the major emission mechanism at low altitudes (below 500 km) and that at higher altitudes the classical bremsstrahlung becomes the major mechanism of emission. However the brightness temperature remains unalterable. Finally, we found that the inverse bremsstrahlung process is not important for the radio emission at these heights. ",
        "introduction": "In order to include the contribution of different ion species, % we use the following (complete) opacity function which resumes the slightly different expressions for the bremsstrahlung opacity reported in the literature \\citep{1979ApJS...40....1K, 1985ARA&A..23..169D,1986rpa..book.....R, 1996ASSL..204.....Z}: \\begin{equation} \\kappa_{ff}^{n}(T,\\nu) = \\sum_\\xi\\frac{2^{5/2}\\sqrt{\\pi}}{3\\sqrt{3}}\\frac{e^6}{c(mk)^{3/2}}\\frac{Z_\\xi^2}{T^{3/2}\\nu^2}n_en_\\xi\\bar{g}_{ff}(Z_\\xi,T,\\nu), \\end{equation} where: $e$ is the electron charge; $c$ is the velocity of light; $m$ the electron mass; $k$ the Boltzmann constant; $Z_\\xi$ is the charge of the ion species $\\xi$; $T$ is the electron temperature; $\\nu$ the frequency; $n_e$ the electron density; $n_\\xi$ is the density of the ion species $\\xi$; and $\\bar{g}_{ff}(Z_\\xi,T,\\nu)$ is the average free-free Gaunt factor which takes into account the initial and final quantum energy states of the total momentum.  The rigorous quantum mechanical expression for the Gaunt factor has been derived by \\cite{1935AnP...415..589S} but the solution is difficult to compute and therefore,  many approximations have been developed in order to make the computation easier (see \\cite{1962RvMP...34..507B} for a review of average Gaunt factor in the bremsstrahlung opacity function).  Following \\cite{1986rpa..book.....R} we % compute the average free-free Gaunt factor for  chromospheric conditions as: \\begin{equation} \\label{laopacidad} \\bar{g}_{ff}(T,Z,\\nu) = \\begin{cases} \\frac{\\sqrt{3}}{\\pi} \\ln \\left[\\frac{4}{\\xi^{5/2}y}x^{1/2}\\right] & x < 1, y < g(x) \\\\ 1 , & x \\le 1 , y > g(x),  y \\le h(x)\\\\ \\left[\\frac{12}{xy}\\right]^{1/2} & x \\le 1,  y > h(x)\\\\ \\frac{\\sqrt{3}}{\\pi}\\ln\\left[\\frac{4}{y\\xi}\\right] & x >1, y < 1\\\\ \\left[\\frac{3}{\\pi y}\\right]^{1/2} & x > 1, y \\ge 1 \\end{cases} \\end{equation} where \\begin{equation} x =  \\frac{kT}{Z^2Ry}, \\end{equation} and $Ry=13.6\\ \\mbox{eV}$ is the Rydberg constant and \\begin{equation} y = \\frac{h\\nu}{kT}, \\end{equation} \\begin{equation} g(x)= \\frac{9999}{10^4}\\left(\\frac{99}{10^2}x+ \\frac{99}{10^3}\\right), \\end{equation} \\begin{equation} h(x)= 1001-1000x. \\end{equation} Under chromospheric conditions it is necessary to include the contribution of the $H^-$ ion in the opacity expression. This  contribution is effective through three mechanisms \\citep{1988A&A...193..189J,1980ZhTFi..50R1847G}:  \\begin{itemize} \\item The Wildt's photo-detachment mechanism \\citep{1939ApJ....90..611W} $$ h\\nu + H^- \\rightarrow H +e^- $$ \\item The neutral interaction \\citep{1996ASSL..204.....Z} $$ h\\nu + e^- + H \\rightarrow H + e^- $$ \\item The bremsstrahlung for $H^-$ or inverse bremsstrahlung \\citep{1980ZhTFi..50R1847G} $$ h\\nu + e^- + H^- \\rightarrow H^- + e^-. $$ \\end{itemize}  In the millimeter - sub-millimeter wavelength range, % the photo-detachment mechanism is not relevant \\citep{1994ApJ...437..879A}, % whereas the neutral interaction is very important. The corresponding opacity function is $$ \\kappa_{ni}^{H} = k_{\\nu}^{tot}P_en_H, $$ where $k_{\\nu}^{tot}$ is the total absorption coefficient published in \\cite{1988A&A...193..189J}, $P_e$ is the electronic pressure and $n_H$ is the Hydrogen density. Finally, in the wavelength range of interest, the inverse bremsstrahlung  opacity is the same than the positive ion case % \\citep[we use bremsstrahlung with $n_i = n_{H^-}$]{1980ZhTFi..50R1847G}.  Therefore, the total opacity used to compute the synthetic spectrum is $$\\kappa_{ff} = \\kappa_{ff}^{n} + \\kappa_{ni}^{H} + \\kappa_{ff}^{H^-}.$$ ",
        "conclusions": "Using different atmospheric models and opacity functions, we have computed the optical depth, $\\tau$, at low  atmospheric altitudes. Our results % show that if we use either, only bremsstrahlung or Celestun (i.e. $H$, $H^-$ and inverse Bremsstrahlung) opacities, the height over the photosphere where the overlying atmosphere becomes optically thick ($\\tau \\approx 1$) remains unalterable. However, if we take into account the local emissivity (Eq. \\ref{eq:eemi}) we show that the height where the main emission is  generated, changes meaningfully.  For example, at 17 GHz the height over the photosphere where the atmosphere starts to become optically thin is around $1500$ km, but the height where $\\tau=1$ is around $2050$ km. There are local emissivity structures below $1500$ km, although they remain invisible to the observer because the atmosphere is optically thick at these heights. When we considered bremsstrahlung opacity only, the local emissivity at 17 GHz changes in a region between $500$ and $850$ km over the photosphere (the atmosphere becomes optically thin). This structure is not present when the Celestun model is considered.   At 212 and 405 GHz the behavior of the local emissivity changes noticeably, since at these frequencies two regions of emissivity are present: A lower one, closer to the photosphere shows deep differences when comparing between both models. The height of the region where the emission is generated is higher by 200 km when the Celestun model is used. The uppermost region is located between $800$ and $1500$ km above the photosphere. In this region there is no difference between the two models. As seen in Figure  \\ref{figure6.ps}, the second layer of emission is below the $\\tau = 1$ height and is similar for both models. Even more, after $1000$ km, the brightness temperature remains unalterable. This is due to the fact that the temperature of the atmospheric model is the upper limit for the brightness temperature, and remains constant between $1000$ and $2000$ km above the photosphere. % the brightness temperature remains unalterable. These facts mask the local emission processes which are  unnoticed by the observer.     We found that the dominant emission mechanism changes with altitude: $H^-$ is dominant below  % $500$ km whereas bremsstrahlung dominates at upper heights, although the final brightness temperature remains unchanged.   Our results show that the inclusion of the $H^-$ and inverse bremsstrahlung mechanisms does not modify the final brightness temperature. But the $H^-$ opacity function has a significant impact in the local emissivity and absorbency and it must be considered when studying micro structures in the low chromosphere. Therefore Celestun code is a valuable tool to perform this kind of studies. The inverse bremsstrahlung has no impact on the brightness temperature, neither on the local emissivity nor on the absorbency processes.  We note that there are more emission mechanisms acting at the low solar atmosphere, as instance, molecular emission has been reported  for CO \\citep{1994Sci...263...64S} and also in the 515.60 to 516.20 nm spectral range  \\citep[$C_2$ and $MgH$, ][]{2002A&A...382L..17F}. % The % emitting layer has been situated below the $1100$ km, making evident that these molecules are present in a thin layer inside the temperature minimum region, between the photosphere and the low chromosphere. However, observations in visible wavelengths, have shown that the molecular emission at the limb brightening region is negligible, corresponding to  0.2\\% of the continuum intensity at the center of the solar disk. Nevertheless, a detailed study of the contribution of the molecular rotational spectrum at these wavelengths is necessary.   Although our simulations are related to quiescent solar situations, the next natural steps will be to investigate and compare with flaring cases. Therefore, such researches could in principle be used to better understand the thermal structure of the flaring solar atmosphere, that is the response of the atmosphere to flare energy release. From the few existing semi-empirical flare models like \\cite{1980ApJ...242..336M} and \\cite{1990ApJ...360..715M} we should grossly expect decreasing radio fluxes as a function of increasing wavelengths in the range from centimeter to far-infrared wavelengths.    We will pay special attention to the density structure of the flaring loops (due to its importance in the computed opacity functions). % It is clear that new simulations, observations and their comparison are needed to better investigate the structure and dynamics of the low layers of the solar atmosphere during flares."
    },
    "1107/1107.2942_arXiv.txt": {
        "abstract": "Cosmological simulations of structure formation follow the collisionless evolution of dark matter starting from a nearly homogeneous field at early times down to the highly clustered configuration at redshift zero. The density field is sampled by a number of particles in number infinitely smaller than those believed to be its actual components and this limits the mass and spatial scales over which we can trust the results of a simulation. Softening of the gravitational force is introduced in collisionless simulations to limit the importance of close encounters between these particles. The scale of softening is generally fixed and chosen as a compromise between the need for high spatial resolution and the need to limit the particle noise. In the scenario of cosmological simulations, where the density field evolves to a highly inhomogeneous state, this compromise results in an appropriate choice only for a certain class of objects, the others being subject to either a biased or a noisy dynamical description. We have implemented adaptive gravitational softening lengths in the cosmological simulation code \\gadgt; the formalism allows the softening scale to vary in space and time according to the density of the environment, at the price of modifying the equation of motion for the particles in order to be consistent with the new dependencies introduced in the system's Lagrangian. We have applied the technique to a number of test cases and to a set of cosmological simulations of structure formation. We conclude that the use of adaptive softening enhances the clustering of particles at small scales, a result visible in the amplitude of the correlation function and in the inner profile of massive objects, thereby anticipating the results expected from much higher resolution simulations. ",
        "introduction": "In collisionless N-body simulations the matter field is represented by a number of point-particles in number considerably reduced with respect to the actual building blocks of the system. These particles have no physical counterpart and should be simply regarded as Monte-Carlo samplings of the probability-density distribution in position and velocity. When during the evolution two of these particles are found at small spatial separations, the gravitational attraction between them can become arbitrarily large and following the encounter properly can turn out to be prohibitively expensive from the computational point of view; plus, this collisional behaviour is an artifact of the granularity of the mass distribution in the simulation and would not manifest in a intrinsically-smooth, collisionless system. A way around this problem, namely the discreetness of the particle representation, is obtained by smoothing the density distribution from a collection of spikes to a continuous field; to this purpose, every particle is assigned a finite volume over which its mass is spread. This translates into the gravitational interaction between the computational particles being ``softened'' at small separations, i.e. being attenuated and prevented to diverge. In this way close encounters are impeded and the simulation can proceed at a regular pace. The gravitational smoothing does not considerably affect the artificial two-body relaxation though, as this is driven by both close and distant encounters and the former cannot be avoided by softening techniques \\citep{chandra42, spitzer71, hernquist_barnes90, theis98, dehnen01, diemand04}. Suppressing relaxation is primarily a condition on the number of sampling particles and depends only weakly on the value of the softening \\citep{power2003}. \\\\ The details of this smoothing procedure can be reduced to the choice of a softening kernel $W(r,h)$, which determines the functional form of the modified density profile of the particle, and of a softening length $h$, which controls  the spatial extent within which the modification applies. In the classic case of ``Plummer'' softening, the gravitational potential of each particle is that of a Plummer sphere whose scale length is given by the value of the softening $h$; this results in the following form for the gravitational force $\\bmath{F}(\\bmath{r})$ between two particles of masses $m_i$ and $m_j$ separated by a distance $\\bmath{r}$: \\begin{equation}\\label{eq:plummer_softening} \\bmath{F}(\\bmath{r}) = -G \\frac{m_i m_j}{(r^2 +h^2)^{3/2}}\\bmath{r}. \\end{equation} An evident shortcoming of this choice is that the gravitational force is modified at all separations; several works (e.g. \\citealt{dyerip93}; \\citealt{dehnen01}) have shown that a better choice is to use kernels with a compact support, meaning that the interaction recovers its Newtonian, original form at separations grater than the softening length. The most commonly used one is the cubic spline of \\cite{ml85}: \\begin{equation}\\label{eq:kernel_ml85} W(r,h) = \\frac{8}{(\\pi h^3)}\\left\\{ \\begin{array}{lcclc} 1-6q^2+6q^3, &\\quad& 0   &\\leq q <& 0.5\\\\ 2(1-q)^3,    &\\quad& 0.5 &\\leq q <& 1\\\\ 0,           &\\quad& 1   &\\leq q \\end{array}\\right. \\end{equation} where $q=r/h$, $r$ being the distance from the centre of the kernel, i.e. from the particle's position. The particles assume the density profile given by Eq.~\\ref{eq:kernel_ml85} and the gravitational potential and force fields are modified accordingly.\\\\ It is evident that the introduction of softening leads to an unphysical modification of the gravitational interaction at interparticle separation below $h$; the inverse square law is replaced by a gentler interaction, whose strength approaches zero in the limit $r \\rightarrow 0$. This results in a systematic misrepresentation of the force at small scales, an effect commonly referred to as \\textsl{bias} \\citep{merritt96,dehnen01}. Ideally one would like to limit this effect to the smallest possible scales by reducing $h$ accordingly; at the same time the softening lengths cannot be made arbitrarily small without running into the discreetness problem previously addressed. It is then clear how the choice of $h$ assumes a critical importance for the reliability of the force computation and indeed a number of studies \\citep{merritt96,romeo98,athanassoula00,dehnen01} have been devoted to the subject.\\\\ In general there is an overall agreement that no such a thing as an ``optimal'' softening length exists when one has to deal with highly inhomogeneous systems, as too big a value would degrade the description of over-dense regions whilst too small a value would enhance collisionality in under-dense environments. Generally, the gravitational softening scale is set to a fixed value, indeed decided as a compromise between the desired spatial resolution and the need to moderate the noise arising from particle-particle interactions. As a matter of fact, in most cases this means fine-tuning the softening parameter to follow properly the dynamics of a certain class of objects \\-- the densest \\-- at a certain time; in a cosmological simulations, where the initially uniform matter density field evolves to form highly differentiated structures, this implies following poorly not only the moderately-dense regions, but also the progenitors of the ``target'', densest objects.\\\\ Ideally it would be desirable to vary the softening scale in space and time according to the density evolution of the different environments: this would remove the aforementioned trade-off between spatial resolution and particle noise, thus allowing to increase the former and reduce the latter depending on the local features of the density field. The problem with a varying softening length though, as will be shown in Section~\\ref{sec:formalism}, is that it introduces an additional dependence on the spatial position $r$ in the definition of the gravitational potential; if this is not accounted for in the derivation of the equation of motion, the energetics of the system can manifest unwanted behaviours. In this context, \\cite{PM07} (hereafter PM07) have proposed a formalism to adapt the gravitational softening lengths in a simulation while retaining conservation of both momentum and energy, which, as just mentioned, would be lost when adopting an adaptive scheme. In their method the softening lengths are computed in the same fashion as for the smoothing length in smoothed particle hydrodynamics (SPH - \\citealt{gm77}; \\citealt{lucy77}), where $h$ varies with the local density $\\rho$ according to $h\\propto \\rho^{-1/3}$. The adoption of this scheme involves a number of modifications in the overall structure of a typical cosmological simulation code, most noticeably the introduction of an additional term in the equations of motion. The technique has already been employed in  simulation codes like the TreePM code of \\cite{BK09} (hereafter BK09) and the TreeSPH \\textsc{EvoL} code of \\cite{merlin}. BK09 apply the method to collisionless simulations of structure formation and find an increase in clustering in over-dense regions accompanied by a somewhat loss in resolution in the under-dense ones; \\cite{merlin} performed standard hydrodynamical tests in order to test the overall performance of their code and found good energy-conservation properties along with an improved behaviour in tests like the isothermal collapse of \\cite{bate97}. Although we are here concerned about particle-based codes, where the gravitational interaction is computed at the particle's location, we note that grid-based codes employing adaptive mesh refinement \\citep{bryan97,truelove98,abel00,knebe01,kravtsov02, teyssier02,oshea04,quilis04} are intrinsically spatially adaptive; these codes perform a recursive refinement of an initial regular grid, thereby increasing their resolution in regions of interest. For similar reasons as those outlined above, conservation of energy is not strictly guaranteed in these codes, as it would not be when using adaptive gravitational softening without changing the equation of motion accordingly.\\\\ In this paper we discuss the implementation of the adaptive softening length formalism by PM07 in the cosmological simulation code \\gadg (\\citealt*{springel01b}; \\citealt{springel05}) and the main effects that it has on dark matter simulations of the large-scale structure of the Universe. Our results qualitatively confirm those of BK09 in over-dense regions, but show no drawbacks in under-dense ones; in particular, we do not register a significant loss of low-mass structures and we considerably improve the results in term of particle clustering when comparing to standard simulations employing fixed softening.\\\\ In Sections~\\ref{sec:formalism} and~\\ref{sec:gadget} we briefly describe the technique and its implementation in \\gadg; in Section~\\ref{sec:tests} we show two of the tests performed to check the implementation and its result in controlled scenarios for which we know the solution, whereas in Sections~\\ref{sec:cosmo} and \\ref{sec:mmII} we apply the technique to cosmological simulations of structure formation. We summarise the results and conclude in Section~\\ref{sec:concl}.  \\section[]{The formalism} \\label{sec:formalism} In this section we will briefly report the gist of the algorithm. We refer the reader to the pivotal work PM07 for a thorough treatment of the problem and for a detailed derivation of the new equation of motion.\\\\ The Lagrangian for a system of $N$ particles interacting only through gravity reads: \\begin{equation}\\label{eq:lagrange} L = \\sum_{i=0}^N m_i\\left(\\frac{1}{2}\\bmath{v}_i^2 - \\Phi_i\\right), \\end{equation} where \\begin{equation}\\label{eq:potential} \\Phi_i(\\bmath{r}_i)= -G \\sum_{j=0}^{N} m_j \\phi \\left(\\left|\\bmath{r}_i - \\bmath{r}_j\\right|, h\\right). \\end{equation} The expression for the potential $\\Phi$ departs from its Newtonian definition due to the introduction of softening and hence of the kernel $\\phi$ and the related length scale $h$. When this quantity has no further dependences and is kept fixed the resulting equation of motion will resemble the Newtonian original, but for the inheritance of this additional scale; its expression is derived by applying the Euler-Lagrange equations to (\\ref{eq:lagrange}) and for particle $i$ it will read: \\begin{equation} m_i\\frac{d\\mathbf{v}_i}{dt} = \\sum_{j=0}^N \\bmath{F}_{ij}(h), \\end{equation} where \\begin{equation} \\bmath{F}_{ij}(h)= -\\nabla \\Phi_{ij} =-G m_j \\phi^\\prime \\left(\\left|\\bmath{r}_i - \\bmath{r}_j\\right|, h\\right)\\frac{\\bmath{r}_i - \\bmath{r}_j}{\\left|\\bmath{r}_i - \\bmath{r}_j\\right|}. \\end{equation} Here $\\phi^\\prime = \\partial \\phi / \\partial\\left|\\bmath{r}_i - \\bmath{r}_j\\right|$ is the force kernel. The expressions for both the potential and the force kernel in the cubic spline case can be found in BK09 (Appendix A) . The force $\\mathbf{F}_{ij}(h)$ reduces to the Newtonian gravitational force for separations $\\left|\\bmath{r}_i - \\bmath{r}_j\\right| > h$ and to the kernel-dependent ``softened'' version otherwise.\\\\ When $h$ is allowed to vary from particle to particle according to the density of the environment a few complications arise. In order to maintain its translational invariance and hence ensure that the resulting dynamics would conserve linear momentum, the Lagrangian needs to be symmetrised. One way to do this is by averaging over the potentials evaluated with the softenings of the interacting pair of particles; the new Lagrangian would read: \\begin{equation}\\label{eq:lagrange_symm} L = \\sum_{i=0}^N \\frac{1}{2}m_i\\bmath{v}_i^2 - \\frac{G}{2}\\sum_{i=0}^N \\sum_{j=0}^N   m_i m_j \\left[ \\frac{\\phi_{ij}(h_i) + \\phi_{ij}(h_j)}{2}\\right], \\end{equation} where $\\phi_{ij}(h_i) \\equiv \\phi \\left(\\left|\\bmath{r}_i - \\bmath{r}_j\\right|, h\\right)$. When deriving the equation of motion from (\\ref{eq:lagrange_symm}), the additional dependence on the position $r$ through $h$ results in an extra term besides the classical inverse square law. The evolution of a particle $i$ subject only to the action of gravity and supplied with a variable softening length will in fact obey: \\begin{eqnarray}\\label{eq:eom} \\frac{d\\bmath{v}_i}{dt} & = & -G \\sum_{j=1}^{N} m_j\\left[ \\frac{\\phi^\\prime_{ij}(h_j) + \\phi^\\prime_{ij}(h_i)}{2}\\right] \\frac{\\bmath{r}_i - \\bmath{r}_j}{\\left|\\bmath{r}_i - \\bmath{r}_j\\right|} \\nonumber \\\\ & & -\\frac{G}{2} \\sum_{j=1}^{N} m_j\\left[ \\frac{\\zeta_i}{\\Omega_i} \\frac{\\partial W_{ij}(h_i)}{\\partial\\bmath{r}_i} + \\frac{\\zeta_j}{\\Omega_j} \\frac{\\partial W_{ij}(h_j)}{\\partial\\bmath{r}_i}\\right], \\end{eqnarray} where \\begin{equation}\\label{eq:zeta} \\zeta_i \\equiv \\frac{\\partial h_i}{\\partial \\rho_i}\\sum_{k=0}^N m_k \\frac{\\partial \\phi_{ik}(h_i)}{\\partial h_i}, \\end{equation} \\begin{equation}\\label{eq:omega} \\Omega_i \\equiv 1 - \\frac{\\partial h_i}{\\partial \\rho_i}\\sum_{k=0}^N m_k \\frac{\\partial W_{ik}(h_i)}{\\partial h_i}. \\end{equation} The additional contribution, which we will refer to as the ``correction term'', is attractive in nature and will therefore act towards increasing the resulting gravitational force. If this term were not accounted for when using adaptive softening, the particles would be evolved according to a law incoherent with the system's Lagrangian, i.e. energy conservation would be lost.  \\section[]{Implementation in Gadget} \\label{sec:gadget} We have implemented the adaptive softening formalism on the latest version of the \\gadg  code, namely \\gadgt. \\gadgt~computes gravitational forces via the TreePM method (\\citealt{xu95}; \\citealt*{bode00}; \\citealt{bagla02}) and hydrodynamical interactions by means of SPH; it differs from the previous versions of the code in that it features a more flexible domain decomposition, something that made it suitable for simulations involving extreme levels of clustering, as the Millennium-II \\citep{mII}.\\\\ Implementing adaptive softening required modifications of the Tree-algorithm and of the timestep criterion along with the introduction of the machinery to compute the softening lengths. Besides the obvious change in the expression for the gravitational acceleration, the opening criterion for the nodes has been modified in order to avoid the presence of smoothed particle-node interaction; in other words, the correction term can be non-zero only within a particle-particle interaction (more details in Sec. 3.2.4 of BK09). The timestep criterion employed by \\gadg for collisionless particles reads: \\begin{equation}\\label{eq:timecriterion} \\Delta t = \\left( \\frac{2\\eta\\epsilon}{a}\\right)^{1/2} \\end{equation} where $\\eta$ is an accuracy parameter, $\\epsilon$ is the Plummer equivalent softening ($h\\simeq 2.8 \\epsilon$, where $h$ is the support of the cubic spline kernel) and $a$ the acceleration of the particle. We kept the criterion itself unchanged and substituted $\\epsilon$ with the individual value of the adaptive softening. \\\\ As to the computation of the softening lengths the method is identical to the one \\gadg uses for setting the smoothing lengths in SPH calculations. Qualitatively, they represent the radius of the sphere centred on the particle and containing a specific number of companion particles, generally referred to as ``neighbours''; formally, this translates into the following relation: \\begin{equation}\\label{eq:what.is.h} \\frac{4\\pi}{3}h^3_i \\sum_{j=0}^N W_{ij}(h_i) = N_{ngbs} \\end{equation} where $N_{ngbs}$ stands for the number of neighbours and $N$ is the number of particles within distance $h_j$ from the target particle $j$. The above equation is solved iteratively with a Newton-Raphson method until the difference between the two sides falls below a certain tolerance threshold.\\\\ Adaptive softening lengths can be activated both when the code works in TreePM mode and when it uses the Tree-only algorithm. In the latter case the softening lengths are left varying without boundaries according to the local features of the particle distribution; in the former case we do instead allow for the presence of a minimum and maximum value for the softening lengths. As explained in Sec.~$3.2.1$ of BK09, the existence of a minimum value is not crucial and only prevents the simulation from becoming overly expensive in terms of computational time; conversely, the upper bound is introduced to ensure that the long-range force (the particle-mesh contribution) is negligible on the scales where softening is important, so that errors arising from the non-modification of the long range force are under control. These bounds are expressed in terms of the splitting scale $r_s$, the scale (generally of order the grid spacing) where the splitting of the potential in a long-range and short-range component is performed; choosing $h_{max} \\simeq \\frac{r_s}{2}$ results in the long-range contribution being below $1\\%$ of the total force at scales where softening is important. Although we always impose a lower limit to the softening length when using the code in its TreePM mode, we did not find the presence of an upper limit to have dramatic consequences on the results, especially when using the adaptive formalism in its full, conservative version.  \\section[]{Tests} \\label{sec:tests} In this section we present some of the tests performed in order to check the correctness of the implementation and explore the general effects of adaptive softening when simulating different physical scenarios. We will initially show the behaviour of the code in simulating simple systems, of well-known properties; the force profile of a Plummer and Hernquist models are investigated, and their temporal evolution in and out of equilibrium; the density profile of a polytrope and the behaviour of its total energy in time is also shown. Most of these examples are present already in PM07 and were specifically chosen to test our implementation.\\\\ In all the numerical simulations presented in this section we adopt units of mass $[M]=1$, length $[R]=1$ and $G=1$. As a result, the energy per unit mass is measured in units of $GM/R$ and time in $(GM/R^3)^{-1/2}$.  \\\\ \\subsection{Evolution of a Plummer sphere} The system considered here consists of a set of $N$ particles distributed according to a ``Plummer'' profile: \\begin{equation} \\rho(r) = \\frac{3GMr_s^2}{4\\pi(r_s^2 + r^2)^{5/2}} \\end{equation} Here the total mass $M$ and the scale radius $r_s$ are set equal to unity. The idea is to evaluate the resulting gravitational force profile and investigate its dependence on the choice of softening. Once this is accomplished we concentrate on the behaviour of the total energy as the system is let evolve in time.\\\\ This test is identical to that presented by PM07 and we refer to section 4.3 of their paper or alternatively to \\citet{aarseth74} for details on the setup of the initial conditions.\\\\ The test was run using different number of particles and both Tree-only and TreePM algorithms for the evaluation of the gravitational force; the results behave as expected in the different cases and here we show only those obtained using the pure Tree method on $N=1000$ particles.\\\\ \\begin{figure} \\includegraphics[width=84mm]{1} \\caption{Average squared errors (ASE) of the force field generated by distributing $N=1000$ particles according to a Plummer profile. The left panel shows the behaviour of the ASE as a function of the choice of (fixed) softening; the right panel shows the results of adaptive softening varying the number of neighbours.} \\label{ase} \\end{figure} Fig.~\\ref{ase} shows the averaged square errors (ASE) of the simulated force field corresponding to different choices of both fixed and adaptive gravitational softening. This quantity measures the deviation of the force experienced by particles at different radii from the analytical value, given by \\begin{equation} f(r) = \\frac{GMr}{(r_s^2 + r^2)^{3/2}} , \\end{equation} and it is defined as \\begin{equation} \\textrm{ASE} = \\frac{1}{f_{max}^2 N} \\sum_{i=1}^N \\left|f_i -f_{exact}(\\bmath{r_i})\\right|^2, \\end{equation} where $f_i$ is the force on particle $i$ and $f_{max}$ is the maximum value of the exact solution. For a discussion on the use of the ASE or related quantities to assess the error on a force profile we refer to PM07, \\citet{merritt96} and \\citet{dehnen01}.\\\\ The results show how less sensitive the representation of the force field is to the choice of $N_{ngbs}$ than to the choice of $\\epsilon$; choosing from $30$ to $500$ neighbours does not shift the global ASE much away from the minimum value corresponding to the ``optimal'' choice of fixed softening ($\\epsilon_{opt} \\approx 0.2$, of the order the mean interparticle separation within the scale radius). As already commented by PM07, the slightly worse behaviour obtained when the new force definition is used and $N_{ngbs} < 100$ can be ascribed to noise-induced gradients in the softening lengths altering the evaluation of the correction term; notwithstanding this, the ASEs for the ``adaptive+correction'' case are remarkably close to the minimum corresponding to the optimal choice of fixed softening throughout the full range of  $N_{ngbs}$. These results compare very well with those of PM07 (see their Fig.~$2$ and the second set of lines from the top). \\begin{figure} \\includegraphics[width=84mm]{2} \\caption{Behaviour of the total energy (kinetic plus potential) of the Plummer sphere as a function of time. The velocities were generated according to the equilibrium distribution function. For a description of the units, see the introduction to Sec.~\\ref{sec:tests}. } \\label{plummer_energy} \\end{figure} \\begin{figure} \\includegraphics[width=84mm]{3} \\caption{Behaviour of the total energy (kinetic plus potential) of the Plummer sphere as a function of time. The equilibrium velocity distribution has been multiplied by a factor $1.2$, leading to an initial global expansion of the system. For a description of the units, see the introduction to Sec.~\\ref{sec:tests}.} \\label{plummer_energy_pert} \\end{figure} Choosing the optimal value for $\\epsilon$ and $N_{ngbs}=60$ we have evolved the system in time and checked the conservation of the total (kinetic plus potential) energy. Fig. \\ref{plummer_energy}  and \\ref{plummer_energy_pert} show the results for two different initial velocity setups, corresponding to a dynamical equilibrium and a perturbed state respectively. In the first case we expect no major evolution of the system  (at least for several relaxation times) whereas in the second an initial overall expansion occurs before a state of equilibrium is reached. We refer again to PM07 and \\citet{aarseth74} for details on the setup of the velocity profiles. The energy fluctuates or increases, reflecting the global changes in the softening lengths, when the adaptive formalism is used without changing the equation of motion accordingly; we register fluctuations of order per cent in the equilibrium set-up and an increase of $\\approx 6\\%$ in the perturbed set-up. Conservation is instead re-established down to time-stepping accuracy as soon as the correct equation of motion is used. Note that the initial total energies are different in the adaptive and fixed softening cases; this is due to the definition and evaluation of the potential energy being different in the two scenarios (see Sec.~\\ref{sec:formalism}).\\\\ \\subsection{Evolution of a Hernquist model} We now perform a similar analysis on a different density distribution; we generate a Monte Carlo realisation of a Hernquist model \\citep{hernquist90}, whose density profile reads: \\begin{equation} \\rho(r) = \\frac{GMr_s}{2\\pi r(r_s + r)^{3}}. \\end{equation} The model is cusped near the origin and provides a more realistic representation of the matter distribution in collapsed objects of cosmological interest, but the steep rise in the density profile makes the distribution more difficult to resolve than the Plummer case addressed in the previous section. We investigate the ASE and the energy evolution for a Hernquist model in equilibrium, sampled with $N=1000$ particles; as in the previous section, the total mass $M$ and the scale radius $r_s$ are set equal to unity. \\begin{figure} \\includegraphics[width=84mm]{4} \\caption{Average squared errors (ASE) of the force field generated by distributing $N=1000$ particles according to a Hernquist profile. The left panel shows the behaviour of the ASE as a function of the choice of (fixed) softening; the right panel shows the results of adaptive softening varying the number of neighbours.} \\label{hern_ase} \\end{figure} Fig.~\\ref{hern_ase} shows the averaged square errors (ASE) of the simulated force field corresponding to different choices of both fixed and adaptive gravitational softening. Again we see how adapting softening provides relatively stable results when varying the number of neighbours from $30$ to $500$; conversely, one needs to choose the fixed softening accurately to avoid a strong misrepresentation of the force profile. We note that the optimal value for $\\epsilon$ is smaller here than for the Plummer sphere ($\\epsilon_{opt} \\approx 0.05$ vs. $\\epsilon_{opt} \\approx 0.2$), a consequence of the different density distribution in the two models: the Hernquist being cusped and the Plummer being flat at the centre. A number of neighbours $N_{ngbs} \\approx 60$ is enough to overcome the noise in the evaluation of the correction term and to provide better results than in the case where plain adaptive softening is used. For any number of neighbours, however, the force profile of the Hernquist model is reproduced better by adapting the gravitational softening. The results are in qualitative agreement with those of PM07 (see their Fig.~$3$ and the second set of lines from the top). \\begin{figure} \\includegraphics[width=84mm]{5} \\caption{Behaviour of the total energy (kinetic plus potential) of the Hernquist model as a function of time. The velocities were generated according to the equilibrium distribution function. For a description of the units, see the introduction to Sec.~\\ref{sec:tests}.} \\label{hern_energy} \\end{figure} As before, we now choose the optimal value for $\\epsilon$ and $N_{ngbs}=60$ to evolve the system in time, checking the behaviour of the total energy. The results are shown in Fig.~\\ref{hern_energy} for the equilibrium model; again, the energy fluctuates ($\\approx 1\\%)$ when the adaptive formalism is used without changing the equation of motion accordingly, whereas conservation is re-established down to time-stepping accuracy as soon as the correct equation is used.\\\\ A more demanding test for energy conservation is to follow the collision of two such particle distributions and check if the adaptive formalism maintains its conservative properties even in this more dynamically active scenario. To this purpose we have simulated the head-on collision of two different realisations of the Hernquist model, where the two systems were placed at a distance $r$ of $150$ scale radii and made approach each other at a speed $v = \\sqrt{(2GM/r)}$. The evolution of the total energy is shown in Fig.~\\ref{hern_energy_merger}; the collision between the two spheres at $t \\simeq 900$ leaves a clear imprint in the total energy, when the softening is allowed to vary and the standard equation of motion is used, while as soon as the correction term is taken into account conservation is recovered.  \\\\ We now want to investigate whether or not the use of adaptive softening helps maintaining the original slope of the density profile in time. To this purpose we let the equilibrium system evolve for several dynamical times and compare the matter distribution among the different models. Fig.~\\ref{hern_profile} shows the cumulative density profile at four different times, marked in the upper left corner. In all the panels the gray line represent the initial profile; the solid lines correspond to the optimal choice for fixed softening ($\\epsilon_{opt} \\approx 0.05$) and to $N_{ngbs} \\approx 60$ for adaptive softening; the dotted lines correspond to $\\epsilon=0.001$ and $N_{ngbs}=30$, whereas the dashed ones to  $\\epsilon=1.0$ and $N_{ngbs}=500$. Again we see how sensitive the results are to the choice of $\\epsilon$ and, correspondingly, how little they depend on $N_{ngbs}$, especially when the correct equation of motion is used. The density profile in the runs with adaptive softening is compatible with the original one at all times (and for all $N_{ngbs}$, when the correction term is added), with a tendency to outperform the results obtained when using fixed softening. \\begin{figure} \\includegraphics[width=84mm]{6} \\caption{Behaviour of the total energy of two colliding Hernquist spheres as a function of time. The collision is head-on, with the two systems at an initial separation $ r = 150\\;r_s$ and approaching each other at a speed $v = \\sqrt{(2GM/r)}$. For a description of the units, see the introduction to Sec.~\\ref{sec:tests}.} \\label{hern_energy_merger} \\end{figure} \\begin{figure*} \\includegraphics{7} \\caption{Cumulative density profile for the $1000$-particles Hernquist model at four different times ($t=200,250,300,350$). The gray lines represent the initial profile at $t=0$; solid lines correspond to the fiducial choice for $\\epsilon (0.05)$ and $N_{ngbs} (60)$; dotted lines correspond to $\\epsilon=0.001$ and $N_{ngbs}=30$, while dashed lines to $\\epsilon=1$ and $N_{ngbs}=500$.} \\label{hern_profile} \\end{figure*} \\subsection{Equilibrium structure of a polytrope} Here we consider a system evolving under the action of both gravity and hydrodynamical forces. A homogeneous gas sphere of initial radius $r_0 = 1$ with equation of state $P=K\\rho^{\\gamma} \\;,\\gamma=5/3$ and initial null internal energy is released to the action of self-gravity and pressure force until hydrostatic equilibrium is reached. The process is speeded up by means of the standard SPH artificial viscosity together with a damping term in the force equation; their effect is to deliberately remove kinetic energy helping the system settling faster to equilibrium. An analytical solution for the internal structure of the system does not exist when $\\gamma=5/3$; the exact radial density profile was then obtained by numerically integrating the corresponding Lane-Emden equation: \\begin{equation}\\label{eq:lanemden} \\frac{\\gamma K}{4\\pi G(\\gamma-1)}\\frac{d^2}{dr^2}(r\\rho^{\\gamma-1}) + r\\rho = 0. \\end{equation} The reference for the setup and realisation of the test is Sec. 4.4 of PM07 and there we refer the reader for additional details on the generation of initial conditions and features of the run.\\\\ \\begin{figure} \\includegraphics[width=84mm]{8} \\caption{Radial density profile of a polytropic sphere in hydrostatic equilibrium. The solid black line represents the numerical solution of the Lane-Emden equation (Eq. \\ref{eq:lanemden}); the points are the results obtained by evolving a $N=1445$-particles system with different choice of softenings.} \\label{poly_prof} \\end{figure} Fig.~\\ref{poly_prof} shows the density profile of the simulated system at $t=40$, plotted against the equilibrium solution given by Eq.~\\ref{eq:lanemden}. The units are as those introduced in the previous section; the time was chosen by following the oscillations in the density at $r=0.2$: when their amplitude reduced of a factor of $20$ with respect to the initial, we assumed equilibrium was reached.\\\\ Three cases are shown, namely fixed softening ($\\epsilon \\approx 1/40$ of the mean interparticle separation, as chosen by PM07) and adaptive softening ($N_{ngbs}=60$) with and without the use of the correction term. In all cases we have used $60$ neighbours for the SPH computations. As found by PM07, the use of adaptive softening provides a better representation of the density profile, especially the inner regions; the higher central densities can be related to a smaller gravitational softening than in the fixed case (at least when the correction term is employed), while the slightly better agreement at large radii comes from much larger softenings reducing the noise of the particle distribution.\\\\ \\begin{figure} \\includegraphics[width=84mm]{9} \\caption{Behaviour of the total energy (kinetic, potential and internal) of the polytropic sphere as a function of time. The system has been perturbed from the equilibrium state at $t=40$ by the induction of radial oscillations. For a description of the units, see the introduction to Sec.~\\ref{sec:tests}. } \\label{poly_ene} \\end{figure} The equilibrium solution just found has then been perturbed by inducing radial oscillations on the sphere. The system was let evolve under gravity and pressure force only, the behaviour of its total energy being now of interest. The results are displayed in Fig. \\ref{poly_ene}. The system oscillates with a period $\\tau \\approx 4$, close to the expected $3.82$ for the fundamental mode of oscillation of such a polytrope. This is seen clearly in the behaviour of the energy when adaptive softening is used and the correction to the equation of motion is neglected; the evaluated potential becomes cyclically deeper and shallower following, respectively, the overall decrease and increase of the softening lengths. Energy conservation is re-established at the level of the runs with fixed softening once the system is evolved according to the appropriate equations.  \\section[]{Performance in a cosmological environment} \\label{sec:cosmo} The problem of the formation and evolution of the large-scale structure of the Universe has no analytical solution and this complicates the analysis when different simulations of the same system are compared; to assess the results and determine which technique is dealing best with the problem is not obvious anymore. Increasing the resolution of a cosmological simulation generally goes along with extending the representation of the initial perturbation field to smaller scales, resulting in new fluctuations entering the initial conditions and therefore, strictly speaking, in a new system. In order to address the problem of structure formation and of how this is dealt with by the different softening formulations we have decided to perform a series of simulations sharing the same initial power spectrum of fluctuations, but differing in the total number of particles (\\citealt*{binney2004}): the same structures form at all the resolution levels, but are more accurately described as more and more particles populate them. If the use of adaptive softening allowed to anticipate the behaviour shown by the standard runs at higher resolution, this would assess the superiority of the method over the usual choice of fixed softening. \\\\ The simulated system here consists of a $100\\; h^{-1}\\mathrm{Mpc}$ side periodic box in a $\\Lambda$CDM cosmology defined by the following choice of parameters: \\begin{eqnarray} & & \\Omega_{\\rm tot} = \\! 1, \\; \\Omega_m = \\!0.3, \\; \\Omega_b=0.04, \\; \\Omega_{\\Lambda}=0.7, \\nonumber \\\\ & & h = 0.7, \\;\\sigma_8=0.9, \\; n_s=1\\,, \\end{eqnarray} where $h$ and $\\sigma_8$ are the values, at redshift zero, of the dimensionless Hubble parameter and rms of the mass fluctuations smoothed on a scale of $8\\; h^{-1}\\mathrm{Mpc}$, whereas $n_s$ is the index of the primordial spectrum of fluctuations. Two sets of simulations have been performed: three runs with $64^3$, $128^3$ and $256^3$ particles using fixed softening and two runs with  $64^3$ and $128^3$ particles using adaptive softening (with and without the correction term). The initial conditions were generated at $z=60$ and differ among them only in the number of particles resolving the same initial fluctuation field. The pre-initial uniform distribution, represented by an equally-spaced grid of particles, is applied a displacement field generated on $64^3$, $128^3$ and $256^3$ grids, depending on the resolution level. The power spectra are anyway identical in all the three cases and they are truncated at the smallest frequency resolved on a $64^3$ grid. The simulations follow only the evolution of the dark matter component of the density field and the resulting mass associated to a single particle varies from $\\approx 32$ to $4$ and $0.5 \\cdot 10^{10}\\;h^{-1}\\mathrm{M}_{\\odot}$ depending on the resolution level of the run. In all the cases we have used the TreePM algorithm varying the grid from $64^3$ to $128^3$ and $256^3$, according to the resolution level. The features of the simulations are summarised in Table~\\ref{table}.\\\\ \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Basic features of the simulations presented in Sec.~\\ref{sec:cosmo}.} \\label{table} \\begin{tabular}{l*{2}{c}r*{2}{c}} \\hline Name                    & $\\epsilon$    & $N_{ngbs}$     & $m_p$                & $k_{min}$           & $k_{max}$        \\\\ & $[h^{-1}kpc]$ &               & $[h^{-1}M_{\\odot}]$  & $[h\\;kpc^{-1}]$      & $[h\\;kpc^{-1}]$   \\\\ \\hline Fixed - $64^3$     \t& $39$  \t& -             & $31.76\\cdot10^{10}$  & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ Fixed - $128^3$    \t& $19.5$        & -             & $3.97\\cdot10^{10}$   & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ Fixed - $256^3$    \t& $10$          & -             & $0.5\\cdot10^{10}$    & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ Adapt - $64^3$     \t& -             & $60 \\pm 0.1$  & $31.76\\cdot10^{10}$  & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ Adapt+corr - $64^3$\t& -             & $60 \\pm 0.1$  & $31.76\\cdot10^{10}$  & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ Adapt - $128^3$    \t& -             & $60 \\pm 0.1$  & $3.97\\cdot10^{10}$   & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ Adapt+corr - $128^3$ \t& -             & $60 \\pm 0.1$  & $3.97\\cdot10^{10}$   & $6.28\\cdot10^{-5}$  &$2\\cdot10^{-3}$  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} The value of fixed softening was chosen to be $1/40$ of the mean interparticle separation in the system, i.e. the smallest advisable value according to the generally accepted criterion that limits this scale to avoid too negative binding energies in close pairs. The choice of the number of neighbours for the adaptive runs corresponds to the minimum number ensuring a robust evaluation of the correction term. To this purpose we have monitored the changes in the two crucial quantities entering the evaluation of the correction term, namely $\\zeta$ and $\\Omega$ (Eq. \\ref{eq:zeta} and \\ref{eq:omega}, respectively), when varying $N_{ngbs}$. The code was run on the $z=0$ snapshot of the $128^3$ simulation employing fixed softening with  $N_{ngbs}$ varying from $16$ to $80$ with a step $\\Delta N_{ngbs}=8$. \\begin{figure} \\includegraphics[width=84mm]{10} \\caption{Number of particle with  $\\zeta$ (Eq.~\\ref{eq:zeta}) in a certain range as a function of $\\zeta$. The code was run on the same clustered distribution of particles using different values for the number of neighbours.} \\label{zeta_distr} \\end{figure} \\begin{figure} \\includegraphics[width=84mm]{11} \\caption{Number of particle with  $\\Omega$ (Eq.~\\ref{eq:omega}) in a certain range as a function of $\\Omega$. The code was run on the same clustered distribution of particles using different values for the number of neighbours.} \\label{omega_distr} \\end{figure} Fig. \\ref{zeta_distr} and \\ref{omega_distr} show the distribution of the two quantities when varying $N_{ngbs}$; we can conclude that a number $N_{ngbs}= 60$ is enough to ensure a converged evaluation of $\\zeta$ and $\\Omega$ and that the resulting correction term would be robustly estimated. We have also run the simulations using adaptive softening and no correction to the equations of motion, using the same number of neighbours as the fully adaptive ones; although the results of the previous section suggest that the introduction of the correction term is crucial to the proper functioning of the adaptive formalism, we will anyway show the behaviour of these simulations for completeness. \\\\ In the following subsections we summarise the main results.  \\subsection{Global behaviour I - Clustering} The level of clustering in the simulations has been first assessed qualitatively by investigating the densities achieved throughout the box. Fig.~\\ref{dens_thres} shows the particles with an associated density greater than $10^5$, $5 \\cdot 10^5$ and $10^6$ times the average density; the results are from the redshift zero snapshots of the $128^3$ and $256^3$ simulations. The densities have been computed in the SPH fashion, using $60$ neighbours for all the runs. \\begin{figure*} \\includegraphics[]{12} \\caption{Particles with an associated density greater than a threshold. Their total number is shown in the upper-right corner. The full $100\\; h^{-1}\\mathrm{Mpc}$ side box is shown and all the particles are plotted; some are not distinguishable from their neighbours due to their proximity, which in some cases can reach $0.5 \\; h^{-1}\\mathrm{kpc}$. The densities were computed as the position-weighted sum of the nearest neighbours, in the usual SPH fashion; the same number of neighbours ($60$) has been used for all the simulations. The results are shown at $z=0$.} \\label{dens_thres} \\end{figure*} As one can see, the runs with adaptive softening reach higher particles densities. Noticeably, the regions where this enhancement of clustering is observed correspond to the high density regions of the $256^3$ simulation; indeed, moving from left to right (i.e. from the fixed softening case to the fully adaptive one) the same areas can be seen to be more and more populated. The numbers in the upper-right corner corresponds to the total number of particles surviving the density threshold; interestingly, the number of these high-density particles in the $256^3$ simulation correspond to a mass of $\\approx 23000,\\; 10000,\\; 300$ particles at the $128^3$ resolution level.\\\\ On a more quantitative level, we have also analysed the clustering by means of the two-point correlation function $\\xi(r)$. This statistics represents the excess probability, compared to a uniform random distribution, of finding pairs of particles at a given spatial separation. \\begin{figure} \\includegraphics[width=84mm]{13} \\caption{Two-point correlation function at $z=0$ for a series of simulations of a $100\\; h^{-1}\\mathrm{Mpc}$ side box in a $\\Lambda$CDM universe; shown is the ratio to the results of the $256^3$ simulation. The runs differ only in the choice of softening and in the total number of particles used to sample the initial perturbation field. The arrows indicate the equivalent-Plummer values of the gravitational softening in the standard runs.} \\label{conv_corr_fct} \\end{figure} Fig.~\\ref{conv_corr_fct} shows the ratio of the correlation functions obtained in the different runs to the result of the  $256^3$ simulation; the curves are shown at $z=0$ and starting from separations of $10 \\; h^{-1}\\mathrm{kpc}$, which correspond roughly to the resolution scale at the $256^3$ level. Other than for some discrepancies at separation of several megaparsecs\\footnote{These discrepancies have no physical meaning and are due to a somewhat less accurate determination of the correlation function at large separations. The method we use is Monte-Carlo based and tuned to provide very solid results at small separations, at the expense of accuracy on large scales. We checked that the differences we register are within the typical uncertainties due to the method when the correlation function is computed at these separations.}, the correlation functions of the three runs with fixed softening overlap almost perfectly down to $100\\; h^{-1}\\mathrm{kpc}$. Below that scale the differences in resolution result in a progressively larger amplitude. Since the behaviour of $\\xi(r)$ at the small-$r$ end is mainly determined by the distribution of particles within individual halos, we could conclude that the internal distribution of particles in a structure becomes denser as the resolution of the simulation is increased. The runs with adaptive softening produce a correlation function which is in full agreement with that of the ``fixed'' run at the same resolution down to roughly $100\\; h^{-1}\\mathrm{kpc}$; at smaller separation the amplitude grows instead larger, approaching the results obtained at higher resolution when using fixed softening. This is particularly evident in the $64^3$-run using adaptive softening and the correction of the equation of motion (``Adapt+corr - $64^3$'', blue curve): not only is the amplitude at small separation ($10$ - $50\\; h^{-1}\\mathrm{kpc}$) higher than the other two $64^3$ runs, but the behaviour at scales  between $50$ and $200\\; h^{-1}\\mathrm{kpc}$ is indistinguishable from the reference $128^3$-run (black, dashed curve). Both the ``adaptive'' runs at the $128^3$ resolution level behave well in reproducing the correlation function of the $256^3$ simulation, the one with correction (purple curve) even slightly exceeding the results for the $256^3$ case on scales below $100\\; h^{-1}\\mathrm{kpc}$.  \\subsection{Global behaviour II - Mass function} The search for structures in the simulations has been carried out with the algorithms \\fof  (\\citealt{fof}) and \\subf (\\citealt{springel01a}). Structures are first identified as collections of $N > N_{min}$ particles separated by mutual distances smaller than some fraction $b$ of the mean interparticle separation. These so-called  ``FOF halos'' are later examined by \\subf for the identification of self-bound substructures and the removal of spurious background particles. In the simulations presented here the halos have been searched using a linking length $b=0.16$ and a minimum threshold of $32$ particles. The internal structure of these candidate objects has then been probed in order to identify local, gravitationally-bound overdensities; those containing at least $20$ particles were addressed to as ``subhalos'' leaving the others as part of the smooth halo component. Finally, particles not gravitationally bound to any substructure of the parent FOF halo were dismissed. Note that \\subf was modified in the unbinding part so that the evaluation of the gravitational potential takes into account the individual softening length of the particles. \\begin{figure} \\includegraphics[width=84mm]{14} \\caption{FOF mass function at $z=0$ for a series of simulations of a $100\\; h^{-1}\\mathrm{Mpc}$ side box in a $\\Lambda$CDM universe. The upper x-axis displays the number of particles in the $128^3$ run corresponding to the masses in the lower axis. The mass functions are truncated at a mass corresponding to $32$ particle masses in every resolution level (i.e. at $\\approx 10^{13},\\; 1.3\\cdot 10^{12},\\; 1.7\\cdot 10^{11} h^{-1}\\mathrm{M}_{\\odot}$ for the $64^3$,$128^3$ and $256^3$ run, respectively). The simulations differ only in the choice of softening and in the total number of particles used to sample the initial perturbation field. } \\label{conv_mfct} \\end{figure} Fig.~\\ref{conv_mfct} shows the mass function of the \\fof halos at $z=0$; the number of objects per logarithmic mass interval is plotted for the different runs down to the mass limit corresponding to $32$ particles. When the correct equation of motion is used, the agreement between the simulations with adaptive and fixed softening at the same resolution level is striking throughout the full mass range; when the correction term is instead ignored, the mass functions drop visibly at low masses, the number of objects containing a number of particles less than of order $N_{ngbs}$ being significantly underestimated. Overall, the runs with adaptive softening and correction underestimate the total number of FOF halos by $\\simeq 3\\%$, as opposed to a $\\simeq 25\\%$ for those without the correction term; if we considered only halos containing more than $100$ particles, these percentages would fall to $\\simeq 1\\%$ and $\\simeq 14\\%$, respectively. Similar results are obtained when considering the ``spherical overdensity'' masses $M_{\\Delta}$\\footnote{These are defined as the masses of the spherical regions centred on the potential minimum of the smooth halo and corresponding to an overdensity of $\\Delta$, typically $\\approx 200$, with respect to either the critical density $\\rho_{crit}$ or to the mean background density $\\rho_{m} = \\Omega_m\\rho_{crit}$. Hereafter, when referring to either the virial radius or the virial mass of an object, we assume the overdensity to be defined with respect to the mean background density.}; this hints to the fact that the global shape of the structures in the simulations is not changed significantly by the adoption of the adaptive softening formalism. \\\\ The upturn in the curve for the $256^3$ simulation at the low-mass end resembles the effect of spurious halos seen in warm dark matter simulations (see, e.g. \\cite{wang07} and their Fig.~$9$); these objects form from the artificial fragmentation of filaments and have masses much smaller than the free-streaming mass of the model being adopted, their origin being entirely due to the specifics of the pre-initial conditions. Having our simulations a truncated power-spectrum, we might in principle be hitting the same problem and have a contribution to our mass functions coming from such spurious structures. This could indeed be the case for the $256^3$ simulation at the very low-mass end, a regime we are not interested in for our comparison. As for the $128^3$ simulations, the problem would affect objects with masses corresponding to a handful of particles that do not survive the $32$-particles limit and therefore do not enter the evaluation of the mass functions. We also specify that this problem does not invalidate the results we displayed in the previous section in terms of correlation functions; the number of particles in the alleged spurious structures in the $256^3$ simulation is insufficient to affect the evaluation of the correlation function at the scales we are interested.  \\subsection{Internal halo properties} The results in terms of correlation function and mass function suggest that the main differences between the runs are likely to manifest in the internal structures of the collapsed objects. We will investigate here what effects the different definitions of softening have in the representation of the most massive halo formed in the simulations. \\begin{figure} \\includegraphics[width=84mm]{15} \\caption{Cumulative mass distribution for the most massive FOF group at $z=0$.} \\label{conv_prof} \\end{figure} Fig.~\\ref{conv_prof} shows the cumulative mass distribution in the halo out to the virial radius. The values differ by fractions of per cent in the various runs and assess around $2.4\\; h^{-1}\\mathrm{Mpc}$. Again it is evident how the use of adaptive softening enhances the clustering of particles anticipating the behaviour of higher resolution simulations. The results of the two runs with adaptive softening and correction (``Adapt+corr - $64^3$'', blue curve; ``Adapt+corr - $128^3$'', purple curve) are particularly worth of notice: the first reproduces the higher-resolution results down to $20\\; h^{-1}\\mathrm{kpc}$, whereas the second is perfectly compatible with the highest resolution run down to less than $5\\; h^{-1}\\mathrm{kpc}$. Another way to interpret this result is in terms of mean inner density as a function of the enclosed number of particles. As noted by \\citet{moore1998} and \\citet{power2003}, obtaining robust results in regions closer to the centre, where the density is higher, demands increasingly large particle numbers. As in Fig. $14$ of \\citet{power2003} we have compared the mean inner density measured at different fractions of the virial radius as a function of the enclosed particle number for all the runs. The results are displayed in Fig.~\\ref{power14}. \\begin{figure} \\includegraphics[width=84mm]{16} \\caption{Mean inner density as a function of the enclosed number of particle at different fractions of the virial radius. The halo under consideration is the same as in Fig. \\ref{conv_prof}. The solid lines separate the regions, at their right, where the average collisional relaxation time ($t_{rel}$) exceeds some fraction of the age of Universe ($t_0$).} \\label{power14} \\end{figure} The $64^3$ simulation using fixed softening provides converged results at most down to $\\simeq 6\\%$ of the virial radius. At smaller radii the results start to diverge from those obtained at higher resolution and cannot be trusted anymore. The two runs with adaptive softening behave instead very well down to $3\\%$ of the virial radius. Similarly, the $128^3$ simulation is reliable down to $\\simeq 2\\%$ of the virial radius when using fixed softening and down to $1\\%$ of the virial radius when adaptive softening is employed. \\citet{power2003} relate the converged side of the enclosed-$\\rho$ vs. enclosed-N plot to the regions where the average collisional relaxation time $t_{rel}$ exceeds some fraction of the age of Universe $t_0$ (between $0.6$ and $1$). In our case, depending on whether we consider the results at $6\\%$ of the virial radius converged or not, we could extend the regions down to the radii where $t_{rel} \\approx 0.4 t_0$. In any case, it is not clear whether the $256^3$ run can be considered converged according to this criterion. It is anyway worth of notice that the $128^3$ adaptive run with correction reproduces perfectly the enclosed density of the $256^3$ at  $1\\%$ of the virial radius.\\\\ Similar results hold when investigating the properties of substructures. Fig.~\\ref{shalos} shows the number of subhalos with mass greater than a certain value; the curves represent the average over the five biggest halos in the simulations and the masses have all been normalised to the virial mass of the host. Due to the poverty of substructures we are not showing the results for the lowest resolution simulations and we are in general limited to qualitative considerations. As can immediately be seen though, the runs with adaptive softening lie closer to the higher resolution simulation than does the run with fixed softening. \\begin{figure} \\includegraphics[width=84mm]{17} \\caption{Average subhalo mass function at $z=0$ for the five most massive halos in the simulations. The subhalo masses have been normalised to the virial mass of the host and vary from $8\\cdot 10^{11}$ to  $4\\cdot 10^{13} h^{-1}\\mathrm{M}_{\\odot}$, corresponding to $\\approx 20$ and $1000$ particles, respectively.} \\label{shalos} \\end{figure}  \\subsection{Comments} All the results shown so far hold at redshift zero. As shown in Fig.~\\ref{soft}, even at this redshift only $\\approx 1 \\%$ of the particles have softening smaller than $1/40$ of the mean interparticle separation and at higher redshifts this behaviour becomes even more extreme. This has the effect of reducing the amplitude of the correlation function more rapidly than in the standard runs when going back in time. A more quantitative representation of this effect will be given in the next section (Fig.~\\ref{mmII_corr}). The mass functions are not substantially affected though, at least not for objects more massive than the $32$-particles limit (Fig.~\\ref{mmII_mfct}).\\\\ As the particle timesteps depends on the value of the gravitational softening (see Eq.~\\ref{eq:timecriterion}), it is natural to expect that they would now span a wider range of values. Fig.~\\ref{timebin} shows that this is the case; plotted are the distributions of particles in time bins\\footnote{In {\\sc GADGET} the simulated timespan is mapped onto the integer interval $[0, 2^{N_{timebins}}]$. This interval is split recursively into timebins, the smallest of which has length $2$. Each time, the code computes individual timesteps for the particles and distributes them in the correspondent bin; the smallest populated bin sets the next timestep for the simulation.} at redshift zero for the three $128^3$ simulations: the ``adaptive'' runs tend to have particles in lower bins than in the ``fixed'' simulation and, at the same time, more particles in the highest bin. This is a consequence of what was just mentioned, namely that only a very small fraction of the particles obtains an adaptive softening smaller than the fiducial choice for the standard simulations at the same resolution level. As finer time bins get populated, the overall number of timesteps increases when adaptive softening is used (along with the number of operations performed within them); at the same time, since the highest bins become more crowded, the average number of active particles \\-- i.e. those that are advanced in a timestep \\-- decreases instead, almost compensating for this overhead. \\begin{figure} \\includegraphics[width=84mm]{18} \\caption{Cumulative distribution of particles in time bins at $z=0$. The simulations were run using $29$ time bins. Smaller numbers corresponds to finer intervals. } \\label{timebin} \\end{figure} Increasing $N_{ngbs}$ has the effect that one intuitively imagines, namely that of increasing the softening associated to the particles; the differences again stand out more evident when looking at the correlation functions, whose amplitude at small separations reduces, whereas the mass functions are confirmed to be considerably less sensitive to variations in $N_{ngbs}$.\\\\ \\begin{figure} \\includegraphics[width=84mm]{19} \\caption{Cumulative distribution of softenings at $z=0$. The softenings are expressed in terms of equivalent-Plummer values ($\\epsilon$). The red (blue) dashed curve represents the value of the gravitational softening in the $128^3$ ($64^3$) ``fixed'' simulations.} \\label{soft} \\end{figure} All the simulations discussed in this section the gravitational softening was prevented to fall below $0.1\\%$ of the TreePM splitting scale $r_s$ (corresponding to $558\\;pc$ for the $64^3$ case and to $279\\;pc$ for the $128^3$ one; in both cases, this translates to $\\approx 1.5\\%$ of the softening scale in the ``fixed'' runs); this implies that the maximum density achievable in the simulation corresponds to roughly $10^9 \\bar{\\rho}$. As pointed out in Sec.~\\ref{sec:gadget}, the presence of a lower bound is not crucial and it is introduced mainly to prevent few particles in over-dense regions from slowing down the simulation. No upper limit to the gravitational softening was set in these simulations; we do not find substantial differences between the runs with and without such a constraint, especially if the correct equation of motion is used. We have also checked both sets against Tree-only runs and found perfectly compatible results.\\\\ Although we both observe an increased level of clustering in massive objects, our global results in terms of particle correlation function and halo mass function somewhat differ from those of BK09; when comparing the results from the adaptive runs to the standard case with fixed softening, BK09 register a deficit in the number of small-mass objects and a slight lowering in the amplitude of the two-point correlation function at small separation, whereas we notice no change in the mass function and instead an overall enhanced level of particle clustering. The simulated background cosmology differ, but this should have no effect in the analysis we are interested here, which is focused on relative differences between the runs more than on absolute results of cosmological interest. The larger linking length we have employed is not at the origin of the discrepancy between the FOF mass functions, nor we think is the different evaluation of the correlation function responsible for the antithetical outcomes.\\\\ The results differ also on the more technical timing level; in our simulations we do not register substantial saving of computing time when adaptive softening is used, but we do notice increasingly better performances when the number of neighbours is increased. BK09 halve the wallclock time of the reference run when using adaptive softening with $32$ neighbours and leave it unchanged when using $48$. If we set an upper limit to the softening of order the splitting scale $r_s$ (as in BK09), we progressively cut the timing when going from $32$ up to $60$ neighbours and leave it essentially unaltered by increasing the number even further; the minimum wallclock time is reached for $N_{ngbs} \\simeq 60$ and it equals that of the simulation with fixed softening. If no upper limit is set for the softening, the wallclock times tend to increase; a minimum is reached for a number of neighbours $\\simeq 40$ and it is around $\\sim 30\\%$ higher than in the standard run.\\\\ We ascribe the responsibility for the different behaviour of the implementation in the various aspects to the respective mother codes. ",
        "conclusions": "\\label{sec:concl} We have implemented adaptive gravitational softening in the cosmological TreePM code \\gadgt. The formalism was introduced first by \\cite{PM07} and features the same technique used by SPH to determine individual softening lengths for each of the simulation particles. The spatial variation of this scale requires a modification of the equation of motion governing the evolution of the particles' trajectories in order to be consistent with the new dependencies introduced in the system's Lagrangian.\\\\  We have applied this technique to several test cases in order to check the behaviour of the total energy when the new equation of motion is used; we then moved to the cosmological scenario and, specifically, to the simulation of the large-scale structure of the Universe, where the evaluation of the effects of adaptive softening is complicated by the lack of an analytical solution to the problem.\\\\ Our main conclusions are: \\begin{itemize} \\item The inclusion of the correction term in the equation of motion is essential to ensure the conservation of total energy. \\item In cosmological simulations of the large-scale structure a number of neighbours $\\approx60$ is needed to obtain a converged estimation of the correction term. \\item With such a choice we show that -- in contrary to previous claims in literature -- the adoption of adaptive softening does not lead to an under-representation of halos at low masses, if the correct equation of motion is used. \\item Using the adaptive scheme effectively increases the dynamical range of cosmological simulations while the computational costs only mildly increase. Especially the amplitude of the two-point correlation function at small scales and the subhalo mass function improve compared to simulations with the same number of particles and fixed gravitational softening, anticipating the results obtained in higher-resolution simulations. \\item The convergence of the inner density profile for the most massive object found in the simulations improves significantly when the adaptive softening and the correct equation of motion is used. \\item When re-simulating a low-resolution version of the Millennium-II simulation \\citep{mII} and comparing the results obtained with fixed and adaptive softening, we notice again perfect agreement in the mass functions at all times and an evolution of the ``adaptive'' correlation function towards the higher amplitude of the Millennium-II's at late times. \\end{itemize} We have so far applied the method to scenarios with equal-mass particle species, either gas-only or dark-matter-only; simulations including both species with different masses would also considerably benefit from the adoption of the scheme: the adaptive behaviour of the resolution scale would allow to follow the collapse of dark matter and particularly gas down to scales currently unachievable in standard simulations at comparable resolution. However, having mixtures of particles with different masses is a non-trivial extension of such scheme and will be investigated in future work. Such future studies will be of particular importance considering that cold gas and star particles in hydrodynamical simulations are likely to clump at lengthscales below those typical chosen for the gravitational softening."
    },
    "1107/1107.0389_arXiv.txt": {
        "abstract": "The theory of cosmological perturbations is extended to spacetimes displaying isotropic expansion but anisotropic curvature. The perturbed Einstein equation and Boltzmann equations for massless and massive particles are derived in a general gauge and a decomposition of perturbations into harmonic modes and moments is proposed. Generalization to the case where anisotropic expansion is also present in the background is discussed. ",
        "introduction": "\\label{intro}  The observed approximate homogeneity and isotropy of the cosmic microwave background (CMB) remains one of the most striking discoveries in cosmology. This approximation holds so well that deviations are adequately characterized by a small perturbation to the temperature $\\Theta\\equiv \\Delta T/T $ as well a polarization of the photon field, expected to be of comparable or smaller order on various scales \\cite{Dodelson:2003ft}. Both of these perturbational quantities are time dependent, inhomogeneous (having a nontrivial spatial dependence) and anisotropic (having a nontrivial dependence upon the angle that an observer see an incoming photon). The influx of data from precision cosmology has allowed for a more searching question to be asked - that of whether these quantities are \\emph{statistically} isotropic. For the case of the temperature perturbation $\\Theta$, fluctuations are statistically isotropic if the following identity holds:  \\begin{eqnarray} \\label{sts} \\left<\\Theta(\\textbf{n}_{(1)}^{i})\\Theta(\\textbf{n}_{(2)}^{i})\\right>= f ( \\textbf{n}_{(1)}^{i}\\textbf{n}_{(2)i}) \\end{eqnarray}  i.e. the above expectation value for the product of temperature fluctuations at two different directions in the sky $\\textbf{n}_{(1)}^{i}$ and $\\textbf{n}_{(2)}^{i}$ is a function of only their projection along one another, the implication being that there is no other spatial direction present in the physics of the generation and evolution of temperature perturbations. The degree to which statistical isotropy of CMB temperature fluctuations is true given the data has been the focus of considerable debate. There exist indications that observations are indeed implying that (\\ref{sts}) is not exact in our universe on large scales, notably there appear to be a number of anomalies that are relatable to the presence of  directions other than  $\\textbf{n}_{(1)}^{i}$ and $\\textbf{n}_{(2)}^{i}$ (see for instance \\cite{Copi:2010na},\\cite{Land:2005ad},\\cite{Hoftuft:2009rq}).  Given these indications one may speculate as to whether the presence of an extra direction in a given CMB anomaly is as a consequence of the presence of fields which impose an actual preferred spatial direction in the universe. This would not be unprecedented in cosmology. For instance, the existence of a field to generate a primordial period of inflation is conventionally taken to explain the existence of a preferred set of clocks on large scales in cosmology \\cite{Brout:1997wg}.  If there does existed a preferred spatial direction, this is expected to be encoded in the gravitational field either at the background or the perturbed level. Consider the following class of background metrics:  \\begin{eqnarray} \\label{back} ds^{2} &=&  a^{2}(t)\\left(-dt^{2}+\\left(\\frac{b(t)}{a(t)}\\right)^{2}dz^{2}+\\gamma_{cd}dx^{c}dx^{d}\\right)\\\\ \\gamma_{cd} &=& \\frac{1}{|K|}\\tilde{\\gamma}_{cd} \\end{eqnarray}  The interpretation of the tensor $\\tilde{\\gamma}_{cd}$ is as follows:  \\begin{itemize} \\item For the case $K>0$, $\\tilde{\\gamma}_{ab}$ is the metric on a 2-sphere with two dimensional Ricci scalar $^{2}R=2$ (i.e unit radius) \\item For the case $K<0$, $\\tilde{\\gamma}_{ab}$ is the metric on the upper sheet of a 2-hyperboloid of two sheets  with two dimensional Ricci scalar $^{2}R=-2$. \\item For the case $K=0$,  the tensor $(1/|K|)\\tilde{\\gamma}_{ab}$ is a two dimensional, flat Euclidean metric (written in cylindrical coordinates). \\end{itemize}  For non-zero values of $a(t)$ and $b(t)$, one may intuitively think of these spaces as an infinite spatial line (the `z direction') with a two-dimensional surface at each point along this line- these surfaces may either a flat 2-space (the case $K=0$), a 2-sphere (the case $K>0$), or the upper sheet of a two dimensional hyperboloid of two sheets ($K<0$).  When $\\frac{db}{dt}\\neq\\frac{da}{dt}$ the expansion rate is anisotropic and hence there exists shear in the congruence of timelike geodesics labelled by proper-time $ \\tau = \\int a dt$. The form of (\\ref{back}) additionally implies that this congruence has vanishing vorticity. Cosmological perturbation theory for the gravitational field coupled to a scalar field for the case with shear and $K=0$ has been developed in detail (\\cite{Pereira:2007yy},\\cite{Pitrou:2008gk},\\cite{Gumrukcuoglu:2007bx}). The case where shear is present and $K\\neq 0$ has also been considered at the level of the background \\cite{Graham:2010hh}. The authors explicitly consider a scenario of primordial anisotropic inflation, followed by isotropic, flat FRW expansion- such expansion corresponds to the choice $K=0$,$a=b$ in (\\ref{back}). The anisotropy is then encoded in the perturbations and not the background. However, we will see in Section \\ref{secstatanis} that this scenario is also straightforwardly addressed in choosing a `2+1' decomposition of the background space as a result of how this allowance for spatial anisotropy is encoded into the perturbation theory.   Alternatively, deviations from Friedmann-Robertson-Walker (FRW) geometry may occur in the absence of shear but in the presence of anisotropic curvature (i.e. $|K| \\neq 0$). This case was discussed \\cite{Mimoso:1993ym} and further developed in \\cite{Coley:1994yt},\\cite{McManus:1994ys}. Crucially, in the absence of shear the metric (\\ref{back}) is conformally static and therefore admits a homogeneous and isotropic solution for the CMB at the background level \\cite{Ehlers:1966ad}. Anisotropic cosmological models which lack the conformally static property typically have fine tuning problems as they they induce anisotropies in the CMB radiation which are known to be consistently regarded as small perturbations around a homogeneous, isotropic solution. Clearly the models are also immune to constraints on directionality of the expansion rate so in the shearless case departures for the predictions of the FRW model begin instead via the anisotropy of the luminosity distance relation which can used to constrain $K$ by comparing with data from supernovae. It was found \\cite{Koivisto:2010dr} that constraints on the quantity $\\Omega_{K}$ (i.e the contribution to the critical density of components with energy density scaling as $a^{-2}$) from the Union SnE dataset \\cite{Kowalski:2008ez} are comparable to those obtained in the FRW model.  The intention of this paper is to develop the theory of cosmological perturbations in the case where the cosmological background belongs to the class (\\ref{back}), is shear free but may posses anisotropic curvature. In Section \\ref{cosper}  we discuss the harmonic decomposition of perturbations and present the perturbed Einstein tensor in a general gauge. Following this we consider perturbations in the matter sectors. Firstly in Section \\ref{sf} we consider perturbations to a putative field supporting the background shear-free condition. Then in Section \\ref{boltmassless} we consider the Boltzmann equation for massless particles, and in doing so propose a decomposition of the angular dependence of the particle distribution function in terms of moments. In Section \\ref{moment} we compare the proposed moment expansion to that conventionally employed in the FRW limit, and in Section \\ref{freest} we discuss the consequences that the mere presence of anisotropic curvature at the background level may have on the evolution of of the photon perturbation. Then, in Section \\ref{boltmassive}, we consider the Boltzmann equation in the case of massive particles. In Section \\ref{polcorr} we present expressions for the angular correlation function for polar temperature perturbations and in Section \\ref{secstatanis} uses these results for comparison to previous results in the literature. The extension of the formalism to the case where shear is present in the background is discussed in Section \\ref{inclshear} and conclusions are presented in Section \\ref{conc}.  Additionally there are a number of appendices: some necessary mathematical background to the calculations is presented in \\ref{appa}; the explicit calculation of the polar angular correlation function is given in \\ref{calpolar}; the form of spacetime gauge transformations and the construction of gauge invariant perturbations is given in \\ref{gaugetran}; comparison to scalar-vector-tensor perturbations in the FRW limit is given in \\ref{compfrw}; and a table of notation is given in \\ref{tableof}. ",
        "conclusions": "\\label{conc}  In this paper we have presented a formalism for the analysis of cosmological perturbations to an anisotropically curved but shearless background. In particular, a harmonic decomposition of spacetime dependent fields has been introduced, as well as a moment decomposition of the dependence these harmonics may in turn have on the direction of particle momentum. One then obtains via the massless and massive particle Boltzmann equations a system of coupled ordinary differential equations in conformal time. These equations are coupled to one another by the presence of adjacent moments of an identical species, possible collisional terms, and to the spacetime via metric source terms. In turn, the Einstein equations describing the evolution of metric perturbations are coupled to the matter fields via their stress energy tensors which can be determined from the fields' collected harmonic moments. Given a set of initial data, these equations are sufficient to describe the evolution of cosmological perturbations.  Of particular interest is the temperature-temperature power spectrum. Explicit expressions for the angular correlation function for polar perturbations have been determined in Section \\ref{polcorr}, and the conditions under which they represent statistical anisotropy have been discussed. In particular, these expressions have been written in a manner which seeks to separate statistical anisotropy which is present via initial conditions (primordial statistical anisotropy) from that which ensues from the evolution of perturbations on an anisotropically curved background and in the presence of perturbations to which fields produce this curvature (which may be termed emergent statistical anisotropy). Indeed, from our formalism we have recovered antecedent results regarding statistical anisotropy of temperature fluctuations induced by an anisotropic primordial power spectrum in an FRW universe (an example of primordial anisotropy, see \\cite{Pitrou:2008gk},\\cite{Gumrukcuoglu:2007bx}) as well as statistical anisotropy induced by a background with anisotropic curvature and anisotropic expansion rate (emergent anisotropy, see \\cite{Graham:2010hh}).  Although in principle the formalism may be used to determine the evolution of cosmological perturbations, there are yet several steps to take in order to be able to accurately assess what kind of observations to expect in a universe with background anisotropic curvature supported by a field such as that discussed in Section \\ref{sf}. Firstly, it is important to provide a careful treatment of primordial power spectra in the presence of anisotropic curvature, where intuitively it would seem that the effects of curvature should be unimportant for the co-moving scales probed today. Secondly, the results in this paper have been presented in a general gauge and the existence of particular gauge choices which may offer considerable simplification has not been determined. Thirdly, polarization in the photon field has not yet been allowed for in calculations, and possible collisional terms (for instance Thompson scattering between electrons and photons) have not yet been explicitly included in the Boltzmann equations. Additionally, it would be desirable to extend the formalism to the case where shear is present in the background. The modifications to the formalism that this should involve have been discussed in Section \\ref{inclshear}.       Acknowledgements: We would like to thank Miguel Quartin, Tomi Koivisto, David Mota, Matthew Mewes, Tim Clifton, Chris Clarkson, Pedro Ferreira, Joe Zuntz, and Mark Wyman for useful discussions."
    },
    "1107/1107.1323_arXiv.txt": {
        "abstract": "{Supernova 1987A (SN\\,1987A) exploded in the Large Magellanic Cloud (LMC). Its proximity and rapid evolution makes it a unique case study of the early phases in the development of a supernova remnant. One particular aspect of interest is the possible formation of dust in SN\\,1987A, as SNe could contribute significantly to the dust seen at high redshifts.} {We explore the properties of SN\\,1987A and its circumburst medium as seen at mm and sub-mm wavelengths, bridging the gap between extant radio and infrared (IR) observations of respectively the synchrotron and dust emission.} {SN\\,1987A was observed with the Australia Telescope Compact Array (ATCA) at 3.2 mm in July 2005, and with the Atacama Pathfinder EXperiment (APEX) at 0.87 mm in May 2007. We present the images and brightness measurements of SN\\,1987A at these wavelengths for the first time.} {SN\\,1987A is detected as an unresolved point source of $11.2\\pm2.0$ mJy at 3.2 mm ($5^{\\prime\\prime}$ beam) and $21\\pm4$ mJy at 0.87 mm ($18^{\\prime\\prime}$ beam). These flux densities are in perfect agreement with extrapolations of the powerlaw radio spectrum and modified-blackbody dust emission, respectively. This places limits on the presence of free--free emission, which is similar to the expected free--free emission from the ionized ejecta from SN\\,1987A. Adjacent, fainter emission is observed at 0.87 mm extending $\\sim0.5^\\prime$ towards the south--west. This could be the impact of the supernova progenitor's wind when it was still a red supergiant upon a dense medium.} {We have established a continuous spectral energy distribution for the emission from SN\\,1987A and its immediate surroundings, linking the IR and radio data. This places limits on the contribution from ionized plasma. Our sub-mm image reveals complexity in the distribution of cold dust surrounding SN\\,1987A, but leaves room for freshly synthesized dust in the SN ejecta.} ",
        "introduction": " ",
        "conclusions": "The first sub-mm (APEX at 0.87 mm) and mm (ATCA at 3.2 mm) observations of SN\\,1987A that we presented here give a first valuable insight into the dust distribution that surrounds the more compact remnant area. These data bridge the gap in the SED between the previously detected far-IR emission from dust grains (at $\\lambda\\leq0.35$ mm) and radio emission from synchrotron radiation (at $\\lambda\\geq0.8$ mm). They do not deviate from extrapolation of the far-IR and radio SEDs, respectively, and we set a limit on the contribution from free--free emission from a plasma, which is close to the estimated contribution from the ionized ejecta from SN\\,1987A. Scheduled high-resolution 3.2-mm observations with the ATCA will improve these constraints.  While the 3.2-mm image ($5^{\\prime\\prime}$ resolution) is point-like the 0.87-mm image ($18^{\\prime\\prime}$ resolution) also reveals emission from cold dust towards the southwest, which may be an interaction region for the RSG progenitor's wind and a dense medium. The (sub)-mm observations presented here provide a valuable guide for preparing high-resolution observations with the Atacama Large Millimeter Array (ALMA), which may settle the question as to the origin of the cold dust that is detected at far-IR (Matsuura et al.\\ 2011) and sub-mm (this work) wavelengths."
    },
    "1107/1107.4992_arXiv.txt": {
        "abstract": "We update the constraints on possible features in the primordial inflationary density perturbation spectrum by using the latest data from the WMAP7 and ACT Cosmic Microwave Background experiments. The inclusion of new data significantly improves the constraints with respect to older work, especially to smaller angular scales. While we found no clear statistical evidence in the data for extensions to the simplest, featureless, inflationary model, models with a step provide a significantly better fit than standard featureless power-law spectra. We show that the possibility of a step in the inflationary potential like the one preferred by current data will soon be tested by the forthcoming temperature and polarization data from the Planck satellite mission. ",
        "introduction": "Current cosmological observations can be explained in terms of the so-called concordance $\\Lambda$CDM model in which the primordial fluctuations are created during an early period of inflationary expansion of the Universe. In particular, the spectrum of anisotropies of the cosmic microwave background (CMB) is in excellent agreement with the inflationary prediction of adiabatic primordial perturbations with a nearly scale-invariant power spectrum \\cite{Komatsu:2010fb,Larson:2010gs,Das:2010ga,Dunkley:2010ge,Hlozek:2011pc}. In its simplest implementation, inflation is driven by the potential energy of a single scalar field, the inflaton, slowly rolling down towards a minimum of its potential; curvature perturbations, that constitute the primordial seeds for structure formation, are originated during the slow roll from quantum fluctuations in the inflaton itself. The scale invariance of the spectrum is directly related to the flatness and smoothness of the inflaton potential, that are necessary to ensure that the slow-roll phase lasts long enough to solve the paradoxes of the Big Bang model.  However, in more general inflationary models, there is the possibility that slow roll is briefly violated. This naturally happens in theories with many interacting scalar fields, as it is the case, for example, in a class of multifield, supergravity-inspired models \\cite{Adams:1996yd, Adams:1997de}, where supersymmetry-breaking phase transitions occur during inflation. These phase transitions correspond to sudden changes in the inflaton effective mass and can be modeled as steps in the inflationary potential. If the transition is very strong, it can stop the inflationary phase as it happens in the usual hybrid inflation scenario; on the contrary, inflation can continue but the inflationary perturbations and thus the shape of the primordial power spectrum are affected. Departures from the standard power-law behaviour can also be caused by changes in the initial conditions due to trans-planckian physics \\cite{Brandenberger:2000wr,Easther:2002xe,Martin:2003kp} or to unusual initial field dynamics \\cite{Burgess:2002ub,Contaldi:2003zv}  A violation of slow-roll will possibly lead to detectable effects on the cosmological observables, or at least to the opportunity to constraint these models by the absence of such effects. In particular, step-like features in the primordial power spectrum have been shown \\cite{Adams:2001vc,Hunt:2004vt} to lead to characteristic localized oscillations in the power spectrum of the primordial curvature perturbation. Such oscillations have been considered as a possible explanation to the ``glitches'' observed by the Wilkinson Microwave Anisotropy Probe (WMAP) in the temperature anisotropy spectrum of the CMB, although the WMAP team notes that these could be caused simply by having neglected beam asymmetry, the gravitational lensing of the CMB, non-gaussianity in the CMB maps and other ``small'' ($\\lesssim 1\\%$) contributions to the covariance matrix. In the following we will assume that these features have indeed a cosmological origin as in the class of extended models described above and we will use CMB data to constrain the phenomenological parameters describing the step in the inflaton potential.  Constraints on oscillation in the primordial perturbation spectrum, as well as best-fit values for the step parameters,  have been previously derived in Refs. \\cite{Peiris:2003ff,Covi:2006ci,Hamann:2007pa,Mortonson:2009qv,Hazra:2010ve}. Here we improve on the previous analyses in several aspects. First, we use more recent CMB data, in particular the WMAP 7-year and the Atacama Cosmology Telescope (ACT) data. This allows us to derive tighter constraints on the parameters; in particular we get an upper limit on the step height (related to the amplitude of oscillations) that is independent on the position of the step itself in the prior range considered. We also find a clear correlation between the position and the height of the step. Secondly, we generate mock data corresponding to the model providing the best-fit to the WMAP data, and use these data to assess the ability of the Planck satellite to detect the presence of oscillations in the primordial spectrum.  The paper is organized as follows: in Section II we describe the evolution of perturbations in interrupted slow roll and the phenomenological model used to describe a step in the inflationary potential. In Section III we discuss the analysis method adopted. In Section IV we present the results and in Section V we derive our conclusions. ",
        "conclusions": "We have considered inflation models with a small-amplitude step-like feature in the inflaton potential. Features of these kind can be due for example to phase transitions occurring during the slow roll in multi-field inflationary models. In these models the primordial perturbation spectrum has the form of a power-law (as in the standard featureless case) with superimposed oscillations, localized in a finite range of scales that basically depends on the position of the step in the potential. We have compared the theoretical predictions of a specific model, i.e., chaotic inflation, and of a more general phenomenological model to the WMAP7 and ACT data, in order to find constraints on the parameter describing the model. We have also studied the possibility of detecting the oscillations with the upcoming Planck data in the case they really exist.  We have found that models with features can improve the fit to the WMAP7 data when the step in the potential is placed in such a way as to produce oscillations in the region $20\\lesssim \\ell \\lesssim 60$, where the WMAP7 data shows some glitches. We found no further evidence for small scales glitches from the recent ACT data, this is fully consistent with the recent analysis of \\cite{Hlozek:2011pc}. We have also found that models with too high a step are excluded by the data. Finally, assuming as a fiducial model the generalized step model that provides the best fit to the WMAP7 data, we have found that the Planck data will allow to measure the parameters of the model with remarkable precision, possibly confirming the presence of glitches in the region $20\\lesssim \\ell \\lesssim 60$."
    },
    "1107/1107.2785_arXiv.txt": {
        "abstract": "The discovery of high-energy ($E>100$~MeV) $\\gamma$ rays from Narrow-Line Seyfert 1 Galaxies ($\\gamma$-NLS1s) has confirmed the presence of powerful relativistic jets in this class of active galactic nuclei (AGN). Although the jet emission is similar to that of blazars and radio galaxies, $\\gamma$-NLS1s have some striking differences: relatively small masses ($10^{6-8}M_{\\odot}$), high accretion rates ($0.1-1$ times the Eddington limit) and are generally hosted by spiral galaxies. It is now possible to study a rather unexplored range of mass and accretion rates of AGN with relativistic jets. Specifically, in this work I present some results obtained by comparing a sample of blazars and $\\gamma$-NLS1s with another sample of Galactic binaries with relativistic jets (stellar mass black holes and neutron stars). ",
        "introduction": "During the latest decades, we have observed powerful relativistic jets from active galactic nuclei (AGN), Galactic binaries (stellar mass black holes, neutron stars and cataclysmic variables), and Gamma Ray Bursts (GRBs). Supersonic jets are also visible in protostellar systems. Despite of the diffusion of jets in almost all types of accreting astrophysical systems, it seemed that galaxies made some preference: early observations of the host galaxies of AGN with relativistic jets displayed only elliptical shapes\\cite{para1,para2,para3,para4}. Only a few counterexamples (i.e. jets in spiral galaxies) were found in the past. The turning point occurred a couple of years ago with the discovery -- by means of the \\emph{Fermi} satellite -- of GeV $\\gamma$ rays from narrow-line Seyfert 1 galaxies ($\\gamma$-NLS1)\\cite{abdo1,abdo2,fosco1}. This class of AGN is composed of sources generally hosted by spiral (barred) galaxies\\cite{crenshaw,deo} and, according to the common paradigm, without jets. Therefore, although some early radio observations had suggested the presence of some jet emission in a handful of cases\\cite{grupe,komossa}, the discovery of GeV $\\gamma$ rays was somehow a surprise. It confirmed that also this class of AGN can develop powerful relativistic jets, like blazars and radio galaxies (see Ref.~\\refcite{fosco2} for a review).  Presently, seven $\\gamma$-NLS1s are known and surely one of them is hosted by a spiral galaxy\\cite{zhou1,anton,fosco2}. Work is in progress to confirm the spiral hosts for the remaining $\\gamma$-NLS1s. The fact that no NLS1 is known to be at high redshift\\footnote{The 13th edition of the V\\'eron-Cetty \\& V\\'eron catalogue of quasars and AGN\\cite{veron} contains only NLS1s with $z\\lesssim 1$. This is also in agreement with the findings that these sources are very young\\cite{mathur}.} suggests that it is very unlikely to find this type of AGN hosted by ellipticals.  The issue of the host is linked to how the mass can affect the generation and the power of the jet and was studied in the past particularly with the use of the radio loudness parameter ($RL=f_{\\rm 5~GHz}/f_{\\rm 440~nm}$), which in turn should indicate the dominance of the synchrotron emission from the jet over the radiation from the accretion disc (e.g. Ref.~\\refcite{para3,para4}). Being the blazars and radio galaxies ($RL>>10$) confined in a mass range of $10^{8-10}M_{\\odot}$, while Seyferts and LINERs ($RL<10$) occupy the range $10^{6-8}M_{\\odot}$, it seemed that the jet needed of a large mass to be generated\\cite{para3,para4}.  The discovery of powerful relativistic jets in $\\gamma$-NLS1 is now filling the region of small masses and high radio loudness (see Ref.~\\refcite{fosco3} for more details). In addition, it opens also new interesting questions on the comparison of jets in AGN and in Galactic binaries. This is the topic of the present work and is complementary to Ref.~\\refcite{fosco3}. ",
        "conclusions": ""
    },
    "1107/1107.5457.txt": {
        "abstract": "A new Monte Carlo method has been developed in order to derive ages of young embedded clusters within massive star forming regions where there is strong differential reddening. After foreground and infrared excess source candidates are removed, each cluster candidate star is individually dereddened. Simulated clusters are constructed using isochrones, an IMF, realistic photometric errors, simulated background field populations and extinction distributions. These synthetic clusters are then dereddened in the same way as the real data, obtained from a deep near infrared survey, and used to derive the ages of 3 embedded clusters. Results were found to be consistent with those determined using spectrophotometric methods. This new method provides way to determine the ages of embedded clusters when only photometric data are available and there is strong differential reddening. ",
        "introduction": "\\label{sec:intro}  Stellar clusters are born embedded within giant molecular clouds (GMCs) and so they are often obscured by many magnitudes of visual extinction. The study of the youngest objects using optical data is often impossible, and therefore infrared data are required to penetrate the natal cloud. Clusters contain groups of stars, covering a large range of stellar masses, confined to a small volume of space. As all cluster members have been formed more or less simultaneously from the same precursor cloud \\citep{lada03}, their study plays an important role in understanding many aspects of stellar evolution. \\\\ $\\indent$Theoretical isochrones, developed using stellar models and the current understanding of pre-main sequence (PMS) evolution, can be used to estimate cluster ages in several different ways. Such an approach is possible as pre-main sequence, main sequence, giant and supergiant stars have very different intrinsic magnitudes and colours, therefore they occupy very different positions on colour-magnitude diagrams (CMDs). The age of a cluster can therefore be estimated by fitting theoretical isochrones to real data in colour-magnitude space.\\\\ $\\indent$The life cycle of a typical cluster may include some or all of the following steps. First a cluster of stars form after the gravitational collapse of a dusty molecular cloud. Gas and dust will fragment throughout the cloud and will eventually form individual stars and systems of stars. In the earliest stage, most stars may still be accreting material and will therefore be obscured by many magnitudes of visual extinction. The cluster at this stage is referred to as an embedded cluster. As the most massive young stars join the main sequence, the subsequent intense UV radiation will begin to ionise the surrounding gas to form an HII region. Through a combination of winds, massive outflows and expanding HII regions, the surrounding natal cloud is slowly dispersed. Once dispersed, all cluster members will have stopped accreting material, may be visible optically and the final mass of all stars in the cluster is set. Such a cluster is referred to as an open cluster.\\\\ $\\indent$Many authors \\citep{bica04,ortolani05,bonatto09} use the theoretical intrinsic magnitudes and colours from computed isochrones to make comparisons with real photometry on apparent colour-magnitude diagrams (CMDs). As open clusters have existed long enough to have dispersed their natal clouds, it can be assumed that all cluster members will suffer from a similar amount of extinction as there is very little, differential reddening between cluster members. \\citet{naylor06} developed an isochrone fitting technique that can be applied in to clusters where a constant amount of reddening can be assumed. More recently work by \\citet{dario10} developed a CMD fitting technique that does allow for some differential reddening between cluster members. They applied their method to pre-main sequence stars in the Magellanic Clouds using Hubble Space Telescope data. However in deeply embedded clusters, due to surrounding gas/dust, the extinction towards cluster members can vary over tens of visual magnitudes. This is particularly true towards massive star forming regions where the extinction inside a cluster can be highly differential. \\\\ $\\indent$The extreme differential reddening, between embedded cluster members, means that it is often impossible to determine the age of an embedded cluster using the traditional techniques of comparing apparent colours and magnitudes with isochrones. However, as dense young clusters probably have nearly all their original stellar population, and low mass members are brighter than at any other time during their main sequence evolution so are easier to count \\citep{lada03}, young clusters offer unique opportunities to study the initial mass function (IMF) and its variation in space and time. Despite the difficulty, for these reasons the study of young embedded clusters are considered very worthwhile. \\\\ $\\indent$Ideally we would spectral type each star in a field of view as the spectral typing of stars can remove the problem of differential reddening. This is because if the spectral type is known, a star can be placed directly onto a Hertzsprung-Russell diagram in order to derive the age. In practice however, due to the observational time required, it is not possible to obtain spectra for the hundreds of cluster members in each of the thousands of clusters contained within the Milky Way. It is however possible to obtain photometry of all of these hundreds of thousands of cluster members, using deep survey data such as the United Kingdom Infrared Telescope Deep Sky Survey (UKIDSS) Galactic Plane Survey \\citep[GPS,][]{lucas08}, and then apply techniques developed using spectroscopy in a systematic fashion to determine parameters such as cluster distance and age.\\\\ $\\indent$In this paper an original approach is used to determine the ages of three embedded clusters, via Monte Carlo simulations, whereby real photometry is dereddened in colour-colour space and then compared in dereddened colour-magnitude space with the photometry of synthetic clusters, generated from theoretical isochrones. In section \\ref{sec:data} we describe the sources of the real data used in this paper. In section \\ref{sec:isochrones_magscolours} we discuss the isochrones used in this paper and the choice of intrinsic colours. In section \\ref{sec:artif_clus} we generate an artificial cluster in order to highlight certain cluster properties and to test the dereddening method used in this paper. In section \\ref{sec:deredProcess} the dereddening process is described. In section \\ref{sec:AgeDet} the age determination process is described and applied to the artificial cluster generated in section \\ref{sec:artif_clus}. The age determination process is extended to three real embedded clusters in section \\ref{sec:real_cluster_ages} and we conclude our results in section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Using a new Monte Carlo method ages have been derived for three different clusters by comparing real dereddened data to artificial data, that can be considered a good approximation to the real data. Stars have been dereddened to an intrinsic locus taking the inevitable problems for early and late type stars into account. Synthetic clusters, created using realistic photometry errors, selection effects and field star contamination, are treated in the same way and therefore account for the same dereddening issues.  \\\\ $\\indent$It was not possible to analyse the most massive stars in the real clusters as they were saturated in UKIDSS. However, as a final test of the age determination method, it is possible to estimate the total luminosity of the real cluster and compare this value with the bolometric luminosities contained in the RMS survey from far-IR fluxes \\citep{mottram11a}. Artificial clusters have been created that are one thousand times the size of each of the three real clusters. Each has the same age as those derived in this paper for the corresponding real clusters. The total luminosity of these artificial clusters has been determined, from the data available in each isochrone, and then scaled down by a factor of one thousand, thereby reducing the effect of any random variation. The calculated total luminosities are 10x10$^4$, 9x10$^4$ and 17x10$^4$ L$_{\\odot}$, and correspond to 2.8, 2.4 and 1.5 times the measured RMS luminosities, for the clusters G042.1268-0.6224, G048.9897-0.2992 and G010.1615-0.3623  respectively. The RMS luminosities should represent a lower limit of the total cluster luminosity. This is because some of the optical/UV radiation will leak out and is not re-radiated in the far-IR. For this reason these values quoted here suggest that a reasonable approximation of the actual cluster luminosities has been performed. \\\\ $\\indent$None of the three clusters show obvious signs of a large age spread within the resolution of the age determination process. However analysis the cluster containing G048.9897-0.2992 had a derived age of 2.0 Myr, this was despite the identification of both an MYSO and UCHII region in the RMS survey. These objects will have typical ages of 0.1 Myr \\citep{mottram11b}, therefore this may suggest multi-epoch star formation rather than continual star formation over an extended period of time. To ascertain this, future work will model a star formation history as opposed to assuming a constant age.\\\\ $\\indent$The age determination method presented in this paper provides a new way to determine the ages of embedded clusters when only photometric data are available. Results were found to be consistent with those determined using more traditional methods, however they have been more reliably determined. Where applicable, the results were also consistent with those derived via spectroscopy. Furthermore, unlike the more tradition photometric age determination methods, this new method can be applied to clusters where the extinction between members is highly differential. This method could be extended to clusters in the southern hemisphere where deeper VISTA data are becoming available \\citep{minniti09}. The deeper data would also yield more reliable results. Finally, tests to this method could be performed using infrared multi-object spectroscopy.  \\appendix"
    },
    "1107/1107.2927_arXiv.txt": {
        "abstract": "In the decade of the dawn of gravitational wave astronomy, the concept of multimessenger astronomy, combining gravitational wave signals to conventional electromagnetic observation, has attracted the attention of the astrophysical community. So far, most of the effort has been focused on ground and space based laser interferometer sources, with little attention devoted to the ongoing and upcoming pulsar timing arrays (PTAs). We argue in this paper that PTA sources, being very massive ($>10^8\\msun$), cosmologically nearby ($z<1$) black hole binaries (MBHBs), are particularly appealing multimessenger carriers. According to current models for massive black hole formation and evolution, the planned Square Kilometer Array (SKA) will observe thousands of such massive systems, being able to individually resolve and locate in the sky several of them (maybe up to a hundred). MBHBs form in galaxy mergers, which are usually accompanied by strong inflows of gas in the center of the merger remnant. By employing a standard model for the evolution of MBHBs in circumbinary discs, with the aid of dedicated numerical simulations, we characterize the gas-binary interplay, identifying possible electromagnetic signatures of the PTA sources. We concentrate our investigation on two particularly promising scenarios in the high energy domain, namely, the detection of X-ray periodic variability and of double broad K$\\alpha$ iron lines. Up to several hundreds of periodic X-ray sources with a flux $>10^{-13}$erg s$^{-1}$cm$^{-2}$ will be in the reach of upcoming X-ray observatories; in the most optimistic case, few of them may be already being observed by the MAXI detector placed on the International Space Station. Double relativistic K$\\alpha$ lines may be observable in a handful of low redshift ($z<0.3$) sources by proposed deep X-ray probes, such as \\textit{Athena}. The exact figures depend on the details of the adopted MBHB population and on the properties of the circumbinary discs, but the existence of a sizable population of sources suitable to multimessenger astronomy is a robust prediction of our investigation. ",
        "introduction": "Within this decade the detection of gravitational waves (GWs) may be a reality, opening a completely new window on the Universe.  While signals coming from compact stars and binaries fall in the observational domain of operating and planned ground based interferometers (such as LIGO, VIRGO, and the proposed Einstein Telescope), massive black hole (MBH) binaries (MBHBs) are expected to be among the primary actors on the upcoming low frequency stage, where the $10^{-4}-10^{-1}$Hz window is going to be probed by gravitational wave interferometry in space (see, e.g., the proposed Laser Interferometer Space Antenna (LISA)). Moving to even lower frequencies, long term monitoring of an array of millisecond pulsars (forming a so called pulsar timing array, PTA) may unveil the characteristic fingerprint left by GWs in the time of arrival of the radio pulses \\citep[e.g.,][]{sazhin78,hellings83}.  The Parkes Pulsar Timing Array \\citep[PPTA,~][]{manchester2008}, the European Pulsar Timing Array \\citep[EPTA,~][]{janssen2008} and the North American Nanohertz Observatory for Gravitational Waves \\citep[NANOGrav,~][]{jenet2009}, joining together in the International Pulsar Timing Array \\citep[IPTA,~][]{hobbs2010}, are already collecting data and improving their sensitivity in the frequency range of $\\sim10^{-9}-10^{-6}$ Hz, and in the next decade the planned Square Kilometer Array \\citep[SKA,~][]{lazio2009} will provide a major leap in sensitivity.  Even though GW observations alone will be an outstanding breakthrough in science, the astrophysical payouts will be greatly amplified by the coincident identification of electromagnetic counterparts. The identification of the host galaxy of a MBH binary merger will (i) improve our understanding of the nature of the galaxy hosting coalescing MBHBs (e.g. galaxy type, colours, morphology, etc.); (ii) help to reconstruct the dynamics of the merging galaxies and of their MBHs; and (iii) offer the possibility of studying accretion phenomena onto systems of known mass and spin \\citep[which, also in PTA observations, can be measured from the GW signal, e.g.][]{vecchio2004,SesanaVecchio10,corbin2010}.  Electromagnetic counterparts to GW events have attracted a lot of attention in the last few years in the context of LISA observations of coalescing MBHBs \\citep[see][and references therein]{schnittman2011}. The final coalescence being a violent event, many possible counterparts related to shocks flares and transients in the surrounding ambient gas have been proposed. However, the bulk of the LISA sources, are likely to cluster at $3<z<8$ \\citep{sesana2007}.  Given the modest sky localization accuracy of the detector \\citep[generally $>1$deg$^2$,][]{lang2008}, the number of candidate hosts in the putative GW errorbox is likely to be very high (order of millions).  Moreover, our incomplete understanding of the observational signatures of such events, together with their intrinsic weakness \\citep[most LISA sources are expected to be MBHBs with masses $\\lesssim10^6\\msun$][]{sesana2007} will make host identification extremely challenging.  On the other hand, very little attention has been devoted to possible counterparts of PTA sources. An obvious advantage with respect to LISA sources is that those are expected to be very massive ($M>10^8\\msun$), cosmologically nearby ($z<1$) systems \\citep{SVV09}. Any characteristic electromagnetic signature would be therefore within the sensitivity range of current and future observation capabilities.  PTA sources are inspiralling MBHBs still far from coalescence (emitting in the $\\sim10^{-9}-10^{-6}$ Hz frequency range), that can be considered, for any practical purpose, stationary during typical observation timescales of decades \\citep{SesanaVecchio10}. Moreover, fairly high eccentricity should be the norm, as shown by MBHB hardening studies in both stellar \\citep{Sesana2010,SGD11,Khan11,Preto11} and gaseous \\citep{Armitage:2005,Jorge09,roedig11} environments. If the binary is surrounded by a circumbinary gaseous disc, eccentricity triggers periodic inflows \\citep{arty96,roedig11}, opening a wide range of possibility for distinctive binary fingerprints that can be observationally identified.  In this paper we study the prospects of conducting multimessenger astronomy with PTA observations of MBHBs. Our aim is to quantify the number of sources possibly detectable both in the GW and in the electromagnetic realm, identifying their signatures, to forecast future identification.  Starting from the Millennium Run \\citep{springel} database{\\footnote{http://www.mpa-garching.mpg.de/galform/virgo/millennium/}}, we construct detailed MBHB populations satisfying all the currently available constraints in terms of local MBH mass function \\citep{marconi2004}, MBH-host relations \\citep{gultekin09}, and MBHB merger rates as inferred from close galaxy pair counts \\cite[e.g.][]{bell2006}.  We assume that after a galaxy merger, cold gas is funneled in the center of the remnant, forming a circumnuclear disc \\citep{mihos96,Mayer07}. The new-formed sub-parsec MBHB excavates a hollow region in the center of the disc (usually referred to as 'gap'), and the dynamics of the system is governed by the mutual MBHB-disc torques. We compute the detailed population of GW emitting sources at each frequency by adopting a simple analytical model for their evolution under the dynamical effect of gaseous drag and GW emission \\cite[see][for details]{KS11}. After identifying the prototypical MBHB observable with PTAs, we simulate its dynamics numerically using a modified version of the smooth particle hydrodynamics (SPH) code GADGET-2 \\citep{Springel05,Jorge09}, to track the fate of the material leaking through the gap and captured by the MBHs. The outcome of the simulation are then used to model analytically the accretion onto the two MBHs and the emitted radiation.  Several distinctive signatures of inspiralling MBHBs have been presented in the literature.  Proposed scenarios range from the radio and the optical, up to the X-rays, including variability of the optical continuum \\citep[e.g.,][]{Haiman2009b}, spectral shifts in the broad line emission \\citep{begelman80,Eracleous11,Tsalmantza11}, peculiar flux ratios between optical/UV broad emission lines \\citep{montuori11}, precessing or x-shaped jets \\citep{liufukun2004,lobanov2005}, periodic outbursts \\citep{valtonen1988}, astrometric measurement of the source motion \\citep{sudou2003}. Yet, the only unambiguous candidate is a double compact radio core at a 7pc projected separation \\cite{Rodriguez2006}.  We identify here a number of possible characteristic signatures and we concentrate our investigation on two particularly promising scenarios in the high energy domain, namely, the detection of X-ray periodic variability and of double broad K$\\alpha$ iron lines. For both scenarios, we quantify the population of potentially observable sources, and the joint detection prospects with future X-ray observatories and PTAs.  The paper is organized as follows. In Section 2 and 3 we model and describe in detail the MBHB population relevant to PTA observations; whereas in Section 4 we investigate the dynamics of the typical source by combining high resolution SPH simulations to an analytical model for the accreted matter. In Section 5 we model the emitted spectrum, identifying possible characteristic signatures, and Section 6 and 7 are devoted to periodic variability and double broad K$\\alpha$ lines, respectively.  We discuss strategy for joint PTA and X-ray detection and draw our conclusions in Section 8. ",
        "conclusions": "The distinctive signatures of PTA-MBHBs highlighted above open interesting scenarios for combining pulsar timing and X-ray observations in the coming years. Even though current PTA efforts (PPTA, EPTA, NANOGrav) may eventually succeed in the detection challenge, multimessenger astronomy prospects are particularly promising in the context of the planned SKA. If a nominal 1ns sensitivity level will be achieved, SKA will make GW astronomy with pulsar timing possible. Several hundreds of signals emitted by the low redshift population of MBHBs will add up to form a confusion noise, on top of which some (possibly a hundred) sources will be individually resolvable and their position in the sky determined.  In this context, X-ray (but also other bands, not considered in this paper) observations may play either a preparatory or a follow-up role. On one hand, the detection of a population of periodic X-ray sources in the present decade, may provide useful information for PTA observations. By the time the SKA will be online, we may have identified a catalogue of systems in the sky from which we expect to detect GW signals, being also able to give indications about the expected frequency. This may substantially facilitate the PTA detection and analysis pipelines, by reducing the search parameter space for specific signals. On the other hand, several resolvable sources with high SNR may be located in the sky with enough accuracy to make follow-up X-ray monitoring possible. To be conservative, according to \\cite{SesanaVecchio10}, typical PTA source sky location accuracy is expected to be in the order of tens of square degrees for a source SNR$=10$ \\citep[but, under specific detection conditions, it can actually be an order of magnitude better, see][]{corbin2010}. According to figure \\ref{resolvable}, typical masses are $>{\\rm few}\\times10^{8}\\msun$ at $z<0.3$. Such systems should be host in massive galaxies with stellar bulges of masses $>10^{11}\\msun$ \\citep{gultekin09}, whose number density is $\\lesssim 10^{-3}$Mpc$^{-3}$ \\citep{bell2003} Considering a sky location error of 10 deg$^2$, the comoving volume enclosed in the error box is $2\\times 10^{-3}$Gpc$^{3}$: maybe up to few thousands candidates falls in the error-box.  Since the presence of a distinctive counterpart relies on the assumption of accretion, active galaxies only should be selected, which may leave us with few dozens of candidates.  Further down selection may be performed by keeping galaxies that show signatures of recent merger activity \\footnote{A detailed study of the statistics of candidate hosts can be found in the independent study by \\cite{tanaka11}, which is complementary to ours in several aspects.}. If just one or at most a few systems are identified, ultra deep X-ray exposure may be used to reveal characteristic double relativistic Fe K$\\alpha$ emission lines.  Such investigation will likely be possible in the next decade with future generation of pulsar timing arrays (in particular with the SKA) and X-ray observatories. In this paper we quantified for the first time the population of expected sources, and their characteristic signatures, and we proposed strategy for coordinating multimessenger observations. We summarize in the following our main results:  \\begin{itemize} \\item About a thousand MBHBs are expected to contribute to the PTA detectable signal in the frequency range $10^{-9}-10^{-6}$Hz, at a level of 1ns or more. Uncertainties in the MBHB population model and in the detailed dynamics of the contributing systems may impact such figures by a factor of a few.  \\item We assumed the standard picture of MBHB migration in circumbinary discs. For popular geometrically thin, optically thick discs described by typical parameters, we find that many of these PTA sources (order of few hundreds, 20\\%-to-50\\% of the total population, for $0.1<\\alpha<0.3$ and $0.1<\\dot{m}<0.3$) are coupled to their circumbinary discs, making a strong case for looking to possible electromagnetic signatures rising by the disc-binary dynamical interplay.  \\item Detailed SPH simulations of the eccentric binary-disc interaction highlight periodic streams of gas leaking from the inner edge of the circumbinary disc, through the low density central gap. The rate at which such streams feed the central binary can be a significant fraction of the Eddington accretion rate of the MBHB (after initial relaxation of the system, we find an average $\\dot{m}\\approx0.4$). The streaming periodicity is extremely sharp, with a variation in the MBHB feeding rate of a factor of two-to-four.  \\item By modelling analytically the dynamics and the emitted spectrum of the inflow-fed minidiscs forming around the two MBHs, we identified several observational features that may be distinctive of an accreting MBHB. We concentrated on the X-ray domain, by separately studying X-ray periodicity and double relativistic Fe K$\\alpha$ lines.  \\item Assuming all merging MBHBs are accreting, we estimate about 100-500 (depending on the detailed property of the circumbinary disc) periodically variable X-ray sources with periods between one and five years. The observed flux on Earth for most of the sources is larger than $10^{-13}$erg s$^{-1}$cm$^{-2}$, within the sensitivity reach of the upcoming X-ray all sky monitor eRosita. However, a lightcurve with at least 15-20 points (eRosita will point each region of the sky at 8 different times only) is required for a statistically significant detection of such a periodicity.  \\item Iron line identification requires deep, targeted observations. We therefore consider as suitable candidate only those PTA sources that are individually resolvable, for which the sky location can be identified to some accuracy. The number of suitable targets is of the order of few tens, depending both on the ability of PTA of resolving sources in the sky, and on the details of the disc-MBHB decoupling and the physical nature of the minidiscs.  \\item Assuming that the accretion flow onto each black hole is capable of generating relativistic Fe K$\\alpha$ emission lines, we have demonstrated that it will in principle be feasible, via high spectral resolution observations with a next generation X-ray observatory (e.g. {\\it Athena}), to identify double K$\\alpha$ line features. Such features can be used to estimate the properties of the black holes, forecasting, in combination with modelling of the PTA signal, the possibility to constrain the space-time around the black hole to unprecedented accuracy.  \\item Note that we assumed {\\it all} MBHB to be surrounded by a circumbinary disc in their late evolutionary stage, prior to merger. Even though this is likely for gas reach merging systems, it might be a too extreme assumption for the massive, low redshift sources relevant to our study. Even assuming that only 10\\% of the systems are surrounded by a circumbinary disc, the number of detectable sources is still interesting: about 10-50 bright, periodic X-ray sources; at least a few individually resolvable PTA sources showing double Fe K$\\alpha$ line profiles.  Even though GW detection of MBHBs remains a challenging task for the present and future astrophysical generations, systematic pulsar timing campaigns are ongoing, and their accuracy and sensitivity will inevitably improve in the coming years. Pulsar timing will therefore provide a safe, open GW window on the low frequency Universe. The combination with electromagnetic observations, such the ones proposed in this paper, will help exploiting GW detection capabilities at their best, providing a lot of information about the population and dynamics of MBHBs.  The present paper is just a first step into the realm of multimessenger astronomy with pulsar timing, hoping that investigators from both the GW and the X-ray/optical/radio communities will take the challenge following our footsteps.  \\end{itemize}"
    },
    "1107/1107.1503_arXiv.txt": {
        "abstract": "{We present a simple (microscopic) model in which bulk viscosity plays a role in explaining the present acceleration of the universe.  The effect of bulk viscosity on the Friedmann equations is to turn the pressure into an ``effective'' pressure containing the bulk viscosity.  For a sufficiently large bulk viscosity, the effective pressure becomes negative and could mimic a dark energy equation of state. Our microscopic model includes self-interacting spin-zero particles (for which the bulk viscosity is known) that are added to the usual energy content of the universe. We study both background equations and linear perturbations in this model.  We show that a dark energy behavior is obtained for reasonable values of the two parameters of the model (i.e. the mass and coupling of the spin-zero particles) and that linear perturbations are well-behaved.  There is no apparent fine tuning involved.  We also discuss the conditions under which hydrodynamics holds, in particular that the spin-zero particles must be in local equilibrium today for viscous effects to be important.} ",
        "introduction": "Experimental evidences --- such as type IA supernovae luminosity~\\cite{Riess_etal_1998,Perlmutter_etal_1998}, Cosmic Microwave Background (CMB) peaks~\\cite{WMAP7_2010}, Baryon Acoustic Oscillations (BAO)~\\cite{Eisenstein_etal_2005} and Large Scale Structures (LSS)~\\cite{Tegmark_etal_2006} --- point toward a universe made of 70\\% ``dark energy'' (for reviews see \\cite{Amendola_Tsujikawa_2010,Frieman_etal_2008}).  This dark energy is responsible for the observed acceleration of type IA supernovae and is thus repulsive (i.e. it has negative pressure).  Explaining this acceleration has become a important challenge for theoretical physics.  There exist many candidate mechanisms/models to explain this acceleration (for reviews see \\cite{Li_Wang_2011,Amendola_Tsujikawa_2010,Caldwell_Kamionkowski_2009,Brax_2009}).  The most famous one is to add a cosmological constant to Einstein's equations.  Although being extremely economical, this solution suffers from severe fine tuning problems. This lead theorists to investigate alternative scenarios, such as modified matter content models (e.g. quintessence~\\cite{Caldwell_etal_1998}), modified gravity models (e.g. scalar-tensor theories~\\cite{Uzan_1999,Amendola_1999,Chiba_1999}) and models relying on the statistical distribution of matter in the universe (e.g. backreaction~\\cite{Rasanen_2003,Kolb_etal_2005,Buchert_2007} or voids~\\cite{Tomita_1999,Celerier_1999}).  The present status is that none of these models offer a satisfactory solution to the dark energy problem because they either suffer from severe fine tuning problems, lead to instabilities, are ruled out or are awaiting further analysis.  In this paper we present an alternative (microscopic) model of dark energy that includes the effects of bulk viscosity.  In an expanding system, relaxation processes associated with bulk viscosity effectively reduce the pressure as compared to the value prescribed by the equation of state.  For a sufficiently large bulk viscosity, the effective pressure becomes negative and could mimic a dark energy behavior.  The idea of having the bulk viscosity drive the acceleration of the universe is mentioned in ref.~\\cite{Padmanabhan_Chitre_1987}.  Unified Dark Matter models take this idea one step further and use a single (bulk) viscous fluid to explain both dark matter and dark energy.  These studies postulate an exotic equation of state for viscous matter of the form $\\zeta = \\zeta_{0}\\rho^{m}$ (where $\\zeta_{0},m$ are parameters)~\\cite{Murphy_1973,Belinskii_Khalatnikov_1975} and investigate its effect on the evolution of the universe~\\cite{Fabris_etal_2005,Colistete_etal_2007,Li_Barrow_2009,Velten_Schwarz_2011}.  It is interesting to note that the Chaplygin gas model~\\cite{Kamenshchik_etal_2001} is a special case of these bulk viscous models.  The conclusion is that it is difficult to explain both dark matter and dark energy using a single viscous fluid~\\cite{Li_Barrow_2009,Amendola_Tsujikawa_2010,Velten_Schwarz_2011}.  The reason for this is that bulk viscosity becomes important at low redshifts (in order to have a negative pressure) and density fluctuations in the viscous fluid are quickly damped.  The resulting decay of the gravitational potential (obeying the Poisson equation) has a dramatic effect on structure formation and leads to an integrated Sachs-Wolfe effect on CMB anisotropies that exceeds observational bounds~\\cite{Li_Barrow_2009,Velten_Schwarz_2011}.  We take a different, more microscopic approach here (for a different microscopic model with bulk viscosity induced by dark matter annihilation, see references~\\cite{Wilson_etal_2006,Mathews_etal_2008}).  Our model consists in adding a new component to the energy content of the universe coming from a self-interacting scalar field.  Thanks to recent progress in thermal field theory, the bulk viscosity of scalar theories has been computed from first principles~\\cite{Jeon_1995}.  The resulting functional form depends on physical properties of the scalar particle (mass and self-coupling) and is very different from previous viscous models.  We study the consequences of this new scalar component (including bulk viscosity) on the evolution of the universe.  We emphasize that our model aims at explaining dark energy only\tand the arguments against Unified Dark Matter models with bulk viscosity do not apply here.  Indeed, the scalar component of our model is added on top of the usual Cold Dark Matter (CDM) component and thus the late damping of fluctuations in the viscous fluid just results in dark energy being very homogeneous today without affecting structure formation or the CMB.  The rest of the paper is organized as follows.  In section~\\ref{sec:Theoretical_concepts} we present some theoretical concepts related to hydrodynamics and bulk viscosity.  We also present our model and the conditions under which it is valid.  Sections~\\ref{sec:Background_evolution} and~\\ref{sec:Linear_perturbations} are dedicated to the study of background equations and linear perturbations respectively.  In each case we derive the modifications to the equations due to bulk viscosity and show the results of simulations done with the Cosmic Linear Anisotropy Solving System\\footnote{\\tt http://class-code.net} ({\\tt CLASS}) code~\\cite{Lesgourgues_2011,Blas_etal_2011}.  The compatibility with observations of our model is presented in section~\\ref{sec:Compatibility_observations}.  We then conclude in Section~\\ref{sec:Conclusion}. ",
        "conclusions": "\\label{sec:Conclusion}  In summary we present a model of cosmological evolution with an additional scalar field.  We show that in a certain range of values for the mass and self-coupling of the scalar field, bulk viscosity plays an important role at late times and mimics a dark energy behavior.  At the background level, the model is compatible with current data on the equation-of-state parameter and predicts a small time variation for this parameter.  At the perturbation level, the bulk viscous model produces the same temperature anisotropies and matter power spectrum as an extended wCDM model, except in the extreme case where the scalar field decouples from the photons at a very late stage and small deviations in the CMB peak heigths are observed.  The model can also account for the extra relativistic degrees of freedom that are marginally preferred by the data.  The Dark Goo model has several features worth mentioning.  For instance, all parameter ranges are bounded, making the model easily falsifiable by future dark energy experiments (note that the upper bound comes from the use of perturbation theory and is not a strict bound in that sense).  The possible values for the model parameters are ``reasonable'' and there is no apparent fine tuning.  On a more philosophical level, it is also reassuring to know that the model has a built-in mechanism (i.e. breakdown of hydrodynamics) that may prevent the universe from accelerating forever (no ``Big Rip'').  The model is also microscopic and the functional form for the bulk viscosity is obtained from first principles.  This allows for a discussion of the validity of the model in terms of microscopic quantities (coupling, mean free path, etc).  This is to be contrasted with other bulk viscous or Chaplygin gas models, where an exotic equation of state with little or no physical justification is postulated.   We stress that the whole model hinges on the validity of hydrodynamics.  The assumptions of hydrodynamics have to be satisfied in order for the concept of bulk viscosity to make sense.  We discuss these issues and provide plausibility arguments for each assumptions, but these arguments can certainly be improved.  In particular, a more fundamental understanding of cavitation would be helpful."
    },
    "1107/1107.1390_arXiv.txt": {
        "abstract": "Motivated by the work of Guillot (2010), we present a semi-analytical formalism for calculating the temperature-pressure profiles in hot Jovian atmospheres which includes the effects of clouds/hazes and collision-induced absorption.  Using the dual-band approximation, we assume that stellar irradiation and thermal emission from the hot Jupiter occur at distinct wavelengths (``shortwave\" versus ``longwave\").  For a purely absorbing cloud/haze, we demonstrate its dual effect of cooling and warming the upper and lower atmosphere, respectively, which modifies, in a non-trivial manner, the condition for whether a temperature inversion is present in the upper atmosphere.  The warming effect becomes more pronounced as the cloud/haze deck resides at greater depths.  If it sits below the shortwave photosphere, the warming effect becomes either more subdued or ceases altogether.  If shortwave scattering is present, its dual effect is to warm and cool the upper and lower atmosphere, respectively, thus counteracting the effects of enhanced longwave absorption by the cloud/haze.  We make a tentative comparison of a 4-parameter model to the temperature-pressure data points inferred from the observations of HD 189733b and estimate that its Bond albedo is approximately 10\\%.  Besides their utility in developing physical intuition, our semi-analytical models are a guide for the parameter space exploration of hot Jovian atmospheres via three-dimensional simulations of atmospheric circulation. ",
        "introduction": "\\label{sect:intro}  The study of extrasolar planets has swiftly transitioned from discovery to characterization.  Following the pioneering detection of the atmosphere of a hot Jupiter by \\cite{char02}, researchers have used the transit method to infer the presence of clouds and hazes in these atmospheres (e.g., \\citealt{pont08,sing11}).  (See \\citealt{fortney10} and references therein for a detailed comparison of the observed atmospheric properties of hot Jupiters to spectral models.)  Such discoveries motivate the improvement of atmospheric models to include the effects of clouds and hazes.  In particular, the simplicity of one- or two-dimensional models allows one to develop physical intuition prior to including new physical effects in (expensive) three-dimensional simulations of atmospheric circulation.  The main purpose of the present study is to generalize the formalism presented in \\cite{hubeny03}, \\cite{hansen08} and \\cite{guillot10} to include several new effects: scattering, collision-induced absorption and additional sources of longwave absorption.  The joint consideration of some of these improvements allows us to understand the effects of clouds/hazes on the temperature-pressure profiles of hot Jovian atmospheres.  The simplicity and versatility of our semi-analytical models allow for an easy comparison to present and future observations aiming to characterize the atmospheres of hot Jupiters.  Absorption and scattering introduce competing effects.  Enhanced longwave absorption by the cloud/haze has the \\textit{dual} effect of cooling and warming the upper and lower atmosphere, respectively.  By contrast, if the cloud/haze scatters in the shortwave, it has the counteracting, dual effect of respectively warming and cooling the upper and lower atmosphere.  All of these effects combine to make the generalization of the condition for the presence or absence of a temperature inversion difficult, as opposed to its crisp one-parameter description in a cloud-free scenario \\citep{hubeny03,hansen08,guillot10}.  As a demonstration of the utility of our models, we tentatively compare our semi-analytical, global-mean temperature-pressure profiles to the data points inferred from the observations of HD 189733b and estimate that its Bond albedo is about 0.1 (Figure \\ref{fig:cloud_fit}).  In \\S\\ref{sect:formalism}, we describe the derivation of a general equation for the temperature-pressure profile in a hot Jupiter atmosphere, which allows for an arbitrary functional form for the longwave opacity $\\kappa_{\\rm L}$.  In \\S\\ref{sect:applications}, we derive more specific forms of the equation assuming various functional forms for $\\kappa_{\\rm L}$.  Examples of temperature-pressure profiles are calculated in \\S\\ref{sect:examples}.  We discuss the implications of our results in \\S\\ref{sect:discussion}.  Table \\ref{tab:symbols} contains a list of the commonly used symbols in this study.  Our main technical results are stated in equations (\\ref{eq:t4_global_cloud0}) and (\\ref{eq:t4_global_cloud}).  Appendices \\ref{append:two_stream}, \\ref{append:deposition}, \\ref{append:identity} and \\ref{append:operators} archive technical information relevant to the formalism presented in \\S\\ref{sect:formalism}. ",
        "conclusions": "\\label{sect:discussion}  \\subsection{Application of models to observations: HD 189733b}  We now make a tentative attempt to compare our models of the global-mean temperature-pressure profile to data points inferred from the published observations of the hot Jupiter HD 189733b.  We focus on HD 189733b as a case study because of the detection of haze in its atmosphere \\citep{pont08,sing09,sing11}.  We make use of the transit spectroscopy, secondary eclipse and photometric phase curve observations of HD 189733b from the \\textit{Hubble} and \\textit{Spitzer Space Telescopes}.  In comparing one- to three-dimensional models, \\cite{fortney10} demonstrated that the temperature-pressure profiles on the exoplanetary limb are good approximations to the true global-mean profile, making it natural to apply our formalism to temperatures derived from transit data.  We extract the hemispherically-averaged brightness temperatures from the phase curves measured using \\textit{Spitzer} at both 8 and 24 $\\mu$m \\citep{knutson07,knutson09}.  For the temperatures at the limb during transit, we use the average value, $(T_{\\rm max}+T_{\\rm min})/2$, as given in Table 3 of \\cite{knutson09} for both the 8 and 24 $\\mu$m temperatures, finding temperatures of 1135$\\pm$52 K and 1102$\\pm$67 K, respectively.  We also use the 3.6, 4.5, 5.8 and 16 $\\mu$m secondary eclipse measurements of \\cite{char08} and \\cite{deming06} with the brightness temperatures given in \\cite{ca11}.  As there is no phase curve information at these wavelengths, we applied a correction of one-half of the average day-night temperature contrast as measured at 8 and 24 $\\mu$m (121$\\pm$35 K) to estimate limb temperatures, finding temperatures of $1518 \\pm 49$ K, $1197 \\pm 57$ K, $1247 \\pm 77$ K and $1217 \\pm 62$ K at 3.6, 4.5, 5.8 and 16 $\\mu$m, respectively.  The assumed pressures for these data points are model dependent and based upon the normalized contribution functions from \\cite{knutson09}.  As such, the absolute pressure scale is uncertain and ultimately tied to the assumed (solar) model abundances and composition, with H$_2$O typically being the dominant opacity source in the near infrared.  Water has been identified in absorption in the \\textit{Spitzer} emission spectra of \\cite{grill08}.  We augment the \\textit{Spitzer} data with the temperatures associated with Rayleigh-scattering haze as detected and confirmed by \\textit{Hubble Space Telescope (HST)} transmission spectra \\citep{pont08,sing09,sing11}.  We use the temperature derived from the \\textit{HST ACS} and \\textit{NICMOS} data in \\cite{sing09} of 1280$\\pm$110 K.  We tie the pressure of the optical haze to that of the \\textit{Spitzer} emission spectra measurements using the differences in altitude between the optical and near-infrared transit spectra as well as the difference from the transit geometry.  Denoting the radius of the hot Jupiter by $R_p$, the 0.75 $\\mu$m optical transit radius is  $7\\times 10^{-4} R_p/R_\\star$ above the 8 $\\mu$m transit radius of \\cite{agol10}, which is approximately 1.7 scale heights at 1300 K or a difference in pressure of a factor of 5.5.  In addition, the slant-to-normal difference of the transit geometry decreases the pressure of the transit measurements by a factor of $\\sqrt{2 \\pi R_p/H} \\approx 50$ at a given wavelength, with $H$ being the pressure scale height \\citep{burrows01,fortney05,hansen08}.  Thus, the haze in the optical at 0.75 $\\mu$m is a factor of about 275 lower in pressure than the 8 $\\mu$m point from the emission spectrum, placing the haze at about the 0.7 mbar pressure level.  This estimated pressure value is in agreement with that of \\cite{les08}, who assumed a haze composition of MgSiO$_3$.  While $P<0.1$ mbar data points do exist (Huitson, Sing et al., in preparation), it is important to note that our formalism is invalid at these low pressures because it does not include the physics of the thermosphere, including hydrodynamic escape (e.g., \\citealt{mc09}).  It is likely that our assumptions of hydrostatic balance and radiative equilibrium break down significantly at these altitudes.  Therefore, we terminate the model comparisons for $P \\le 0.1$ mbar.  The initial task is to cut down on the number of parameters used in our models.  We estimate that $T_{\\rm irr} \\approx 1697$ K $(1-{\\cal A})^{1/4}$ \\citep{bouchy05} and $g \\approx 21.88$ m s$^{-2}$ \\citep{torres08}.  Since the data points are for $P < 1$ bar, they set no constraint on the temperatures at depth and thus are unaffected by the value of $T_{\\rm int}$ adopted --- for simplicity, we set $T_{\\rm int}=0$ K.  Again for simplicity, we set $\\kappa_{\\rm c_0}=0$ cm$^2$ g$^{-1}$, namely that any extra longwave absorption by the cloud/haze present may be assimilated into the longwave opacity function $\\kappa_{\\rm L}$.  Therefore, we are left with a 4-parameter model to be compared to 7 data points.  The assumption of a uniform scattering opacity is not unreasonable given the fact that \\cite{sing11} detect a transmission spectrum of HD 189733b consistent with a haze layer, producing Rayleigh scattering, at $\\sim 1$ mbar, spanning about 8 pressure scale heights (albeit at somewhat higher altitudes than where we are making our model comparisons).  \\begin{figure} \\begin{center} \\includegraphics[width=0.48\\columnwidth]{f8a.eps} \\includegraphics[width=0.48\\columnwidth]{f8b.eps} \\end{center} \\vspace{-0.2in} \\caption{Comparison of global-mean temperature-pressure profiles to data points inferred from observations of HD 189733b (see text for details and caveats).  The observed system parameters allow us to set $T_{\\rm irr}=1697$ K, while we set $T_{\\rm int}=0$ K and $\\kappa_{\\rm c_0}=0$ cm$^2$ g$^{-1}$ for simplicity. The left panel fixes the shortwave opacity and longwave opacity normalization and illustrates the effects due to the variation of the other two parameters.  The right panel fixes $\\epsilon$ and $\\xi$ and varies the opacities.  In the left panel, the horizontal, dash-triple-dot line shows the approximate position of the longwave photosphere, while the hatched, yellow region indicates the approximate location of the shortwave photon deposition depth (corresponding to the pair of $\\xi$ values adopted).  We terminate all of the models at $P=0.1$ mbar as our formalism does not treat the physics of the thermosphere.  For comparison, we include the temperature-pressure profile from Madhusudhan \\& Seager (2009), labelled ``MS09\", which was obtained via an abundance and temperature retrieval method.} \\label{fig:cloud_fit} \\end{figure}  We begin with a $\\gamma_0=1.5$ model with some scattering present ($\\xi=0.7, {\\cal A} \\approx 0.1$) and without including the effect of collision-induced absorption ($\\epsilon=1$), using equation (\\ref{eq:t4_global_cloud0}), as shown in the left panel of Figure \\ref{fig:cloud_fit}.  It is clear that while such a model is roughly consistent with the increase in temperature from $\\sim 0.1$ bar to $\\sim 1$ mbar, it fails to produce the increase in temperature with pressure at $P \\gtrsim 0.1$ bar.  By contrast, a model with no scattering ($\\xi=1, {\\cal A}=0$) and which includes collision-induced absorption ($\\epsilon=3500$) is roughly consistent with both of these trends, but does not quite pass through the data points.  As we retain $\\epsilon=3500$ and allow scattering to be present ($\\xi=0.7, {\\cal A}\\approx 0.1$), the lower atmosphere becomes cooler, thus allowing the model to be consistent with the data points.  Too much scattering ($\\xi=0.5, {\\cal A} \\approx 0.2$) results in a temperature-pressure profile which is cooler than indicated by the data points.  It is reasonable to ask if our model comparisons are sensitive to the assumed values of the opacities.  In the right panel of Figure \\ref{fig:cloud_fit}, we retain $\\gamma_0=1.5$ but calculate two more models where we reduce and increase both of the shortwave and longwave opacities by a constant factor of 2.  It is clear that even with $\\gamma_0$ kept constant, the ``turn\" in the temperature-pressure profile --- i.e., first decreasing, then increasing, temperature with increasing altitude --- is sensitive to the values of the opacities.  Lower opacities result in the ``turn\" residing at higher pressures.  The temperatures at depth ($P \\gtrsim 1$ bar) vary sensitively when the opacities are varied due to the changing depths at which the stellar photon energy is being deposited (see equation [\\ref{eq:deposition}]).  For comparison, we include in Figure \\ref{fig:cloud_fit} the temperature-pressure profile for HD 189733b from \\cite{madhu09}, which was obtained via an abundance and temperature retrieval method applied to infrared transit observations from both the \\textit{HST} and the \\textit{Spitzer Space Telescope}.  Given the uncertainties associated with both theory and observation, the profile is broadly consistent with our models for $P \\gtrsim 0.1$ bar but does not include the temperature inversion from $\\sim 0.1$ bar to $\\sim 1$ mbar.  This is unsurprising because the data point at about 1 mbar was not included in the original analysis of \\cite{madhu09}.  One may find the status of HD 189733b as the prototypical example of an inversion-less hot Jupiter to be at odds with the increase in temperature from $\\sim 0.1$ bar to $\\sim 1$ mbar and our interpretation of it using a $\\gamma_0 > 1$ model.  This dilemma is resolved if one realizes that there are two types of temperature inversions: deep- versus meso-atmospheric inversions.  The former refers to temperature inversions probed by near infrared water lines, typically at $\\sim 0.1$--1 bar: at these pressure levels, there is general consensus that HD 189733b and HD 209458b are the prototypes of hot Jupiters without and with temperature inversions in their atmospheres, respectively \\citep{burrows07,char08,fortney08,madhu09}.  The latter refers to temperature inversions from $\\sim 0.1$ bar to $\\sim 1$ mbar, as probed by the \\textit{HST} transmission spectra of \\cite{pont08} and \\cite{sing09,sing11}.  \\textit{Collectively, our models make a crude prediction for the Bond albedo of HD 189733b: ${\\cal A} \\approx 0.1$.}  The Bond albedo is not an observable quantity; rather, it is the geometric albedo which is measured, as has been done for HD 209458b by \\cite{rowe08}.  (See \\citealt{burrows08} for the implications of albedo measurements on modelling.)  The conversion from the geometric to the Bond albedo requires integrating over both frequency and phase angle.  For example, \\cite{bs09} estimate that the \\textit{spherical albedo}\\footnote{The Bond albedo is obtained by integrating the spherical albedo over frequency.} ${\\cal A}_s$ and the geometric albedo ${\\cal A}_g$ are related by ${\\cal A}_g \\approx 0.8 {\\cal A}_s$ in their models of exoplanetary atmospheres with Rayleigh scattering.  The reader is again referred to \\S3.4 of \\cite{seager10} for a detailed discussion of exoplanetary albedos.  \\subsection{Semi-analytical models as a guide for three-dimensional simulations}  The ability of our models to predict the temperature at depth ($T_\\infty$) for a given set of parameters provides a useful guide to three-dimensional simulations of atmospheric circulation which utilize dual-band, two-stream radiative transfer \\citep{hfp11}.  Such simulations require the specification of the initial temperature and velocity fields.  The simplest assumption is to initiate them from a state of windless isothermality, which then requires the specification of a constant, initial temperature $T_{\\rm init}$.  Computational efficiency is optimized when the initial temperature field is as close to radiative equilibrium as possible.  Physically, we expect the temperature to approach radiative equilibrium at depth ($P \\gtrsim 10$ bar).  Therefore, selecting $T_{\\rm init}=T_\\infty$ is a good choice.  As the simulation proceeds, the upper atmosphere ($P \\lesssim 10$ bar) adjusts itself to temperatures consistent with dynamical-radiative equilibrium, while the lower atmosphere ($P \\gtrsim 10$ bar) remains in radiative equilibrium.  A state of quasi-equilibrium is reached when there is dynamical-radiative equilibrium within the upper atmosphere, while the lower atmosphere ceases to transfer significant amounts of energy upwards.  Semi-analytical models such as the ones presented in this study alleviate the need to perform a tedious parameter search for the appropriate value of $T_{\\rm init}$ to adopt and therefore increase the efficiency of utilizing general circulation models to study the atmospheres of hot Jupiters.  \\subsection{Future work}  A useful extension of our models will be to consider the dependence of the radius, at a given wavelength, of a hot Jupiter on various quantities, \\begin{equation} R = R\\left( T_\\infty, T_{\\rm int}, g \\right). \\label{eq:radius} \\end{equation} The temperature at depth $T_\\infty$ is a reasonable proxy for the atmospheric effects due to the variation of absorption and scattering.  Our formalism can then be used to estimate a value for $T_\\infty$, which will provide a lower boundary condition for models of hot Jovian interiors.  While the elucidation of equation (\\ref{eq:radius}) is beyond the scope of the present study, such an exercise will be useful for interpreting the inflated radii of some hot Jupiters (e.g., \\citealt{laughlin11}).  Specifically, combining our atmospheric models with those of hot Jovian interiors will allow a connection to the observed (transit) radius, thereby allowing us to rule out models which over-inflate a given hot Jupiter.  \\subsection{Summary}  The salient points of our study may be summarized as follows: \\begin{itemize}  \\item We have presented a semi-analytical model for the global-mean temperature-pressure profile of a hot Jovian atmosphere which requires the specification of 7 parameters: the intrinsic heat flux (via $T_{\\rm int}$), the flux of stellar irradiation (via $T_{\\rm irr}$ or $T_{\\rm eq}$), the surface gravity of the exoplanet ($g$), the strength of shortwave scattering ($\\xi$ or via ${\\cal A}$), the shortwave opacity ($\\kappa_{\\rm S}$) and the longwave opacity (via $\\kappa_0$ and $\\epsilon$).  Allowing for an extra source of longwave absortion due to the presence of a uniform cloud/haze layer adds another parameter ($\\kappa_{\\rm c_0}$), while a purely absorbing Gaussian cloud/haze deck adds another 2 parameters to the system ($\\kappa_{\\rm c_0}$, $\\Delta_{\\rm c}$ and $P_{\\rm c}$; but sets $\\xi=1$).  \\item If a cloud/haze layer contributes an extra source of longwave opacity, then the main effect is to cool and heat the upper and lower atmosphere, respectively, analogous to the greenhouse effect on Earth.  \\item If a cloud/haze layer causes scattering in the shortwave, then the main effect is to heat and cool the upper and lower atmosphere, respectively, thus counteracting its absorbing effect in the longwave --- the anti-greenhouse effect.  Scattering shifts both the shortwave photosphere and the photon deposition depth to higher altitudes.  \\item Besides their utility in developing physical intuition, our semi-analytical models provide a guide to three-dimensional simulations of atmospheric circulation.  In particular, they pin down the temperature at depth for a given set of parameters, which serves as the initial condition for the temperature field.  \\item We have made a tentative comparison of 4-parameter models to the temperature-pressure data points inferred from the \\textit{Hubble} and \\textit{Spitzer} observations of HD 189733b and estimate that its Bond albedo is approximately 0.1.  Future observations of HD 189733b will corroborate or refute some of our models.  \\item The simplicity and versatility of our semi-analytical models allow for an easy comparison to observations.  Observers wishing to use our models should refer to equation (\\ref{eq:t4_global_cloud0}) and set $T_{\\rm int}=0$ K and $\\gamma^{-1}_{\\rm c}=0$ as a starting point, thereby distilling the model down to having only 4 parameters.  \\end{itemize}  \\begin{table*} \\centering \\caption{Table of commonly used symbols} \\label{tab:symbols} \\begin{tabular}{lcc} \\hline\\hline \\multicolumn{1}{c}{Symbol} & \\multicolumn{1}{c}{Meaning} & \\multicolumn{1}{c}{Units}\\\\ \\hline \\vspace{2pt} $\\mu = \\cos\\theta$ & cosine of latitude & --- \\\\ $m$ & column mass per unit area of atmosphere & g cm$^{-2}$ \\\\ $m_0$ & column mass per unit area at bottom of model atmosphere & g cm$^{-2}$ \\\\ $P$ & vertical pressure & bar or dyne cm$^{-2}$ \\\\ $P_0$ & pressure at bottom of model atmosphere & bar or dyne cm$^{-2}$ \\\\ $P_{\\rm c}$ & pressure level where cloud/haze deck is located & bar or dyne cm$^{-2}$ \\\\ $\\xi$ & shortwave scattering parameter (ratio of absorption to total opacity) & --- \\\\ $\\epsilon$ & factor for enhancement of longwave optical depth (at $P=P_0$) due to collision-induced absorption & --- \\\\ $\\kappa_{\\rm S}$ & shortwave opacity & cm$^2$ g$^{-1}$ \\\\ $\\kappa_{\\rm L}$ & longwave opacity & cm$^2$ g$^{-1}$ \\\\ $\\kappa_0$ & longwave opacity normalization & cm$^2$ g$^{-1}$ \\\\ $\\kappa_{\\rm c}$ & additional longwave opacity due to cloud/haze deck & cm$^2$ g$^{-1}$ \\\\ $\\kappa_{\\rm c_0}$ & longwave opacity normalization of cloud/haze deck & cm$^2$ g$^{-1}$ \\\\ $\\gamma \\equiv \\kappa_{\\rm S}/\\kappa_{\\rm L}$ & ratio of shortwave to longwave opacities & --- \\\\ $\\gamma_0 \\equiv \\kappa_{\\rm S}/\\kappa_0$ & ratio of shortwave opacity to longwave opacity normalization & --- \\\\ $\\gamma_{\\rm c} \\equiv \\kappa_{\\rm S}/\\kappa_{\\rm c_0}$ & ratio of shortwave opacity to longwave opacity normalization due to cloud/haze & --- \\\\ $\\tau_{\\rm L} = \\int \\kappa_{\\rm L} ~dm$ & longwave optical depth & --- \\\\ $\\tau_{\\rm S} = \\kappa_{\\rm S} m/\\xi$ & shortwave optical depth & --- \\\\ $\\tau_0 \\equiv \\kappa_0 m_0$ & longwave optical depth normalization in absence of collision-induced absorption & --- \\\\ $\\tau \\equiv \\kappa_{\\rm L} m$ & --- & --- \\\\ $\\tau_{\\rm c}$ & additional longwave optical depth due to cloud/haze deck & --- \\\\ $\\Delta_{\\rm c}$ & thickness parameter of cloud/haze deck & --- \\\\ $T_{\\rm int}$ & blackbody-equivalent temperature associated with internal heat flux & K \\\\ $T_{\\rm irr}$ & irradiation temperature & K \\\\ $T_{\\rm eq} = T_{\\rm irr}/\\sqrt{2}$ & equilibrium temperature of hot Jupiter & K \\\\ $\\bar{T}$ & global-mean temperature of atmosphere & K \\\\ $T_\\infty$ & temperature at depth & K \\\\ \\hline \\hline \\end{tabular}\\\\ Note: the term ``opacity\" refers only to absorption opacities. \\end{table*}  \\vspace{0.2in} \\noindent \\textit{K.H. acknowledges support from the Zwicky Prize Fellowship and the Star and Planet Formation Group at ETH Z\\\"{u}rich, and benefited from stimulating interactions with the Astrophysics Group at Exeter University.  W.H. acknowledges support by the European Research Council under the European Community's 7th Framework Programme (FP7/2007--2013 Grant Agreement no. 247060).  F.P. acknowledges support by a STFC Advanced Fellowship.  We thank Isabelle Baraffe, Tristan Guillot, Esther Buenzli, Adam Burrows, Hans Martin Schmid, Sascha Quanz, Giovanna Tinetti and Nikku Madhusudhan for useful conversations.}  \\appendix"
    },
    "1107/1107.4867_arXiv.txt": {
        "abstract": "In this paper we combine the WMAP7 with lookback time and Chandra gas fraction data to constrain the main cosmological parameters and the equation of state for the dark energy. We find that the lookback time is a good measurement that can improve the determination of the equation of state for the dark energy with regard to other external data sets. We conclude that larger lookback time data set will further improve our determination of the cosmological parameters. ",
        "introduction": "The seventh-year Wilkinson Microwave Anisotropy Probe (WMAP) data rigorously test the standard cosmological model placing constraints on its basic parameters. The WMAP measurements alone are not enough to break the degeneracy among some cosmological parameters or to place constraints on non-standard cosmological models. For example, measurements of cosmic microwave background (CMB) power spectrum alone do not strongly constrain the curvature of the universe characterized by the energy density parameter $\\Omega_k$. One needs to complement the CMB data with the luminosity or angular diameter distances measurements in order to constrain $\\Omega_k$ because the astrophysical distances depend also on the expansion history of the universe \\cite{komatsu-wmap7}.  There are conclusive evidences that the universe is in a state of accelerated expansion. The Hubble diagram of Type Ia supernovae (SNeIa, \\cite{Hicken09, schaf09, guimaraes09}), combined with  CMB anisotropy measurements (\\cite{komatsu-wmap7, dunkley09}), baryon acoustic oscillations (BAO) from the galaxy distribution data (\\cite{percivaletal07, samushia09a, gaztanaga09, wang09}), and galaxy cluster gas mass fraction measurements (\\cite{allen08, samushia08, ettori09}) support the idea that we live in a spatially-flat universe where nonrelativistic matter make almost 30\\% of the critical energy density while the rest is an unknown component called dark energy, with negative effective pressure being responsible for the present phase of accelerated  expansion of the universe \\cite{samushia09}.  Constrains on dark energy density parameter $\\Omega_\\Lambda$ and on its equation of state $w$ (the ratio of pressure to energy density), can explain the nature of the repulsive force causing the acceleration of the universe. One possible explanation for this unknown component is an energy density constant in time and uniform in space. Such a cosmological constant ($\\Lambda$) was originally postulated by Einstein to explain a static universe, later rejected when the expansion of the Universe was first detected and presently reinstated to account for the dark energy. Still, the computed value of $\\Lambda$ is expected to be $10^{120}$ larger than the observed one. Another cosmological scenario considers that dark energy is a dynamical scalar field with a time varying equation of state. An alternative explanation of the accelerating expansion of the Universe is that general relativity or the standard cosmological model is incorrect.  However, ground and space observations can not discriminate among different dark energy scenarios as the correct explanation of the observed accelerating universe: a cosmological constant, a dynamical scalar field or a modification of general relativity  \\cite{DETF, peebles02}.  The WMAP7-year data combined with other astrophysical measurements \\cite{komatsu-wmap7} place constraints on the dark energy. Assuming a flat universe ($\\Omega_k \\sim 0$), an accurate determination of the Hubble expansion rate ($H_0$) helps in improving the limit of the equation of state of the dark energy \\cite{spergel03, hu05}. In the paper of Komatsu et. all. \\cite{komatsu-wmap7} from the joint analysis of WMAP7+BAO+$H_0$ in the case of a time independent equation of state, a value of $w = -1.10 \\pm 0.14$ at $68\\%$ CL was obtained. Furthermore, adding high-z supernova data to their analysis a more stringent limit was obtained, $w = -0.98 \\pm 0.053$ at $68\\%$ CL. However, this later result does not take into account the systematic errors in supernovae, which are comparable with the statistical errors \\cite{kessler09, hicken09}. Also, combining the cluster abundance and 5-year WMAP data, Vikhlinin et al. \\cite{vikhlinin09} found that for a flat universe $w = -1.08 \\pm 0.18$ at $68\\%$ CL. Furthermore, adding BAO \\cite{eisenstein05} and the supernova data \\cite{davis07}, they found $w = -0.991 \\pm 0.09$ at $68\\%$ CL.  In this paper, we perform a joint analysis of the CMB-WMAP7 data, constraints on Hubble expansion rate inferred from the age of astrophysical objects using lookback time method (LBT) and measurements of the gas mass fraction of relaxed clusters from Chandra X-ray observatory(Chandra). Our aim is to investigate for several parameters of a given cosmological model which combination of this three data sets puts better constrains.  We choose the lookback time method because it has the advantage of using the ages of distant objects which are independent of each other so we can avoid biases present in techniques that use distances of primary or secondary indicators in the cosmic distance ladder method \\cite{samushia09}.  Moreover, we use Chandra measurements because they currently provides one of the best constraints on $\\Omega_m$ and have the advantage of being remarkably simple and robust in terms of its underlying assumptions \\cite{allen08}.  The paper is structured as follows: in section 2 we briefly present the lookback time method, in section 3 Chandra gas fraction experiment, in section 4 we describe the statistical analysis and present the data sets and finally, in the last sections, the results and conclusions. ",
        "conclusions": "The aim of our paper was to analyze how the LBT and Chandra data sets combined with WMAP7 measurements can improve the determination of the equation of state for the dark energy and also of the other parameters of the standard cosmological model. We choose the LBT external data set because it contains information independent of each other and we can avoid biases present in the cosmic distance ladder method and the Chandra data set because it provides one of the best constrains on the matter density parameter $\\Omega_m$.  Our contribution consist in implementing LBT and Chandra new modules into the public Monte Carlo Markov Chain package. We run the modified CosmoMC package for the following combinations of data sets: WMAP7, WMAP7+LBT, WMAP7+Chandra, WMAP7+LBT+Chandra. We found that, the physical baryonic density parameter, the scalar spectral index, the equation of state for the dark energy, the Hubble expansion rate, the age of the universe and the amplitude at the pivot scale are best constrain using WMAP7+Chandra data sets. All the other considered parameters, the physical dark matter, dark energy and matter density parameters and the plane $\\Omega_{m} - \\Omega_{\\Lambda}$ are best constrain when we combine WMAP7+LBT+Chandra.  Moreover, we conclude that the looking back time is a trustful measurement and in the future a larger LBT data set can further improve the determination of parameters we investigate.  Acknowledgements. This work is supported by the ESA/PECS Contract \"Scientific exploitation of Planck-LFI data\" and by CNCSIS Contract 539/2009."
    },
    "1107/1107.0991_arXiv.txt": {
        "abstract": "We have analysed the long- and short-term time dependence of the pulse arrival times and the pulse detection rates for eight Rotating Radio Transient (RRAT) sources from the Parkes Multi--beam Pulsar Survey (PMPS). We find significant periodicities in the individual pulse arrival times from six RRATs. These periodicities range from $\\sim$30 minutes to 2100 days and from one to 16 independent (i.e. non--harmonically related) periodicities are detected for each RRAT. In addition, we find that pulse emission is a random (i.e. Poisson) process on short (hour--long) time scales but that most of the objects exhibit longer term (months--years) non--random behaviour. We find that PSRs J1819$-$1458 and J1317$-$5759 emit more doublets (two consecutive pulses) and triplets (three consecutive pulses) than is expected in random pulse distributions. No evidence for such an excess is found for the other RRATs. There are several different models for RRAT emission depending on both extrinsic and intrinsic factors which are consistent with these properties. ",
        "introduction": "\\label{intro} Rotating Radio Transients (RRATs) are neutron stars which were discovered only through their isolated pulses \\citep{mll+06}. Some, however, have later been detectable through periodicity searches. The average intervals between detected pulses range from a few minutes to a few hours and pulses have durations between 2 and 30 ms. Thus far, $\\sim$50 RRATs have been identified \\citep[][]{hrk+08,dcm+09,klk+09,bb10,kkl+11,bbj+11}, including the original 11 from \\citet{mll+06}. Periods ranging from 0.1 to 8 seconds have been measured for 29 of these sources.  Period derivatives have been measured for 14, allowing inference of spin--down properties such as characteristic ages and surface dipole magnetic fields \\citep[][]{mlk+09,lmk09,kkl+11}. % The periods and magnetic fields of RRATs are larger than those of normal pulsars, but the distributions of other spin--down properties such as spin--down energy loss rate and characteristic age are similar \\citep{mlk+09}. Despite this overall trend, the properties of individual RRATs vary considerably. Four RRATs, including PSRs J1826$-$1419 and J1913+1330, have spin--down properties consistent with the bulk of the  normal radio pulsar population and two others, PSRs J1317$-$5759 and J1444$-$6026, have properties similar to normal, older pulsars. Four others, PSRs J1652$-$4406, J1707$-$4417, J1807$-$2557, and J1840$-$1419, lie just above the radio `death--line' \\citep[e.g.][]{cr93,zgd07}. However, some have more unusual spin--down properties. PSRs J0847$-$4316, J1846$-$0257, and J1854$+$0306,  lie in an empty region of $P-\\dot{P}$ space between the normal radio pulsars and isolated neutron stars (XINS) and PSR J1819$-$1458 has a high magnetic field of $5\\times10^{13}$~G.   Because of the difficulties in detecting these sporadic objects, the total Galactic population of RRATs likely outnumbers that of normal radio pulsars \\citep{mll+06}, though it is possible that the populations are evolutionarily related \\citep{kk+08}. Several ideas have been presented about the nature of the emission from these objects. It could be similar to that responsible for the `giant pulses' observed from some pulsars \\citep[e.g.][]{kbm+06}.  It could also be that the sporadic emission is related to the fact that these objects are near the radio `death--line' \\citep[e.g.][]{cr93,zgd07} and/or are examples of extreme nulling \\citep[e.g.][]{rr09}. The phenomenon has also been attributed to the presence of a circumstellar asteroid belt \\citep{li06,cs08} or a radiation belt as seen in planetary magnetospheres \\citep{lm07}. Or, perhaps, some are transient X-ray magnetars \\citep[e.g.][]{wkg+05}. Another idea is that their properties lie at the extreme end of the population of normal radio pulsars. Weltevrede et al. (2006) \\nocite{wsrw06} show that  PSR B0656$+$14, a nearby middle--aged pulsar which emits pulses with energies many times its mean pulse energy, would be discovered as a RRAT source if it were farther away. RRATs may also be considered as an extreme case of mode changing \\citep[see, e.g.,][]{wmj+06} where the on state is less than or about one pulse period. Furthermore, \\citet{lhk+10} have recently shown that many pulsars exhibit a two--state phenomenon in which varying pulse profile shapes are correlated with variations in spin--down rates and implied changes in magnetospheric particle density. These changes are quasi--periodic, with timescales ranging from one month to many years. It could be that the RRATs are similar two--state systems, in which the profile changes are so dramatic to make them undetectable in the more common state.  Determining the time variability and/or periodicity of the RRAT pulses is therefore an important diagnostic of the RRAT emission mechanism. While the pulse profile, and, in most cases, pulse intensity changes of  \\citet{lhk+10} are quasi--periodic, the pulse intensity distributions of normal pulsars and giant--pulsing pulsars are believed to be random over time. On the other hand, nulling pulsars in general show on and off timescales of more than one consecutive pulse, indicating largely non--random distributions \\citep[see, e.g.,][]{rr09}. Radio emitting neutron stars often show  transient spin--down phenomena as well. For instance, glitches, or sudden increases in the spin frequency, have been observed from young pulsars and one RRAT ( PSR J1819$-$1458). One of the glitches from PSR J1819$-$1458 was accompanied by a 3.5$\\sigma$ increase in the pulse detection rate \\citep{lmk09}. Radiative events do not normally accompany the glitches of normal radio pulsars, but are quite common for magnetars \\citep{dkg08}. This, along with the high magnetic field of  PSR J1819$-$1458, hints at a relationship with magnetars and also provides additional motivation to examine the pulse rate variations with time for all RRATs.  We search for periodicities and quantify the randomness of the detected RRAT pulses in several different ways. We first search for periodicities in the pulse arrival times on minutes--year long time scales and pulse detection rates on month--year long time scales using a Lomb--Scargle analysis. We then quantify the randomness of the RRAT pulse arrival times using Kolmogorov--Smirnov tests on seconds--year long time scales. The observations are described in Section~2, the methods and results in Section~3, and the conclusions and plans for future work in Section~4. ",
        "conclusions": "We have shown that  six of the RRATs have periodicities of significance greater than 99$\\%$ in the pulse arrival times. The periods of the most significant peak range from 1.4 hours (for PSR J1754$-$30) to 2102 days (for PSR J1819$-$1458). No significant periodicities were detected upon randomizing the time series, showing that these periodicities are real. We do not find any relationship between the number and significance of detected periodicities and spin--down properties such as period or characteristic age. It is possible that some of the periodic  behavior is due to refractive scintillation. However, the number and wide range of timescales of the periodicities are impossible to explain with refractive scintillation alone.   The shorter timescale periodicities in pulse arrival times are similar to typically observed nulling timescales, which range from minutes to days \\citep{wmj+06}. Explanations for pulsar nulling include an empty sight--line passing through the sub--beam structure \\citep{dr01}, a reversal of the emission direction \\citep{mg+06}, pulsar emission ceasing temporarily due to intermittent failure of pair production (Zhang et al. 2007), an asteroid belt of material \\citep{cs08} or changes in magnetospheric currents \\citep{lhk+10,wmj+06}. Any combination of these could explain the extreme pulse-to-pulse variability of the RRATs and also the longer term periodicities.   The significant periodicities found in the RRAT pulse arrival times may suggest a relationship with pulsars whose spin--down rates and pulse shapes un- dergo periodic variations, with implied changes in magnetospheric particle density \\citep{lhk+10}. The prototype of this source class is B1931+24 (Kramer et al. 2006). The asteroid belt model of \\citep{cs08} attributes the 40-day on/off timescale of this pulsar to an asteroid with eccentric 40-day orbit. It may be that the pulse--to--pulse variability of the RRATs is due to a similar process happening on very short timescales. However, because the on states of the RRATs are so short, it is impossible to measure period derivatives during the on and off states. It could also be that a similar process is causing the longer term periodicities in pulse arrival times.  If we apply the model of \\citep{cs08} to the RRATs,  multiple asteroids of an asteroid belt could be responsible for the observed periodicities in the arrival times. This is consistent with the large root mean square timing residuals which range from 1.1 ms (for PSR J1913$+$1330) to 11.2 ms (for PSR J0847$-$4316) of these RRATs as an earth--sized asteroid would induce residuals of the order of 1 ms \\citep{cs08}. However it is possible that much of these large residuals are due to pulse--to--pulse jitter, indicating that the true asteroid mass cannot be determined by the residuals only.  The periodic fluctuations in pulse arrival times could also be due to non--radial oscillations which drive different emission modes, as often seen in white dwarf stars \\citep{rr11}. The fundamental oscillation periods for neutron stars are expected to be on the order of milliseconds \\citep{rg92} to seconds \\citep{mh88} for g--modes. These are far too short to explain the multiple periodicities seen, but it is possible that we are observing the beat frequency between a non--radial oscillation period and that from a longer timescale process like those observed in \\citet{lhk+10}.   We also searched for periodicities in the daily pulse detection rates. Six of the RRATs do not show any significant periodicities in their daily pulse detection rates over timescales of months to years. The exceptions are PSR J1819$-$1458, which exhibits a 1260 day period with a significance of 97\\%, and PSR J1754$-$30, which exibits a 994 day period with significance of 95\\%.  We detect a peak of higher significance only 0.6\\% and  1.0\\% of the time in 1000 trials in which the rates were randomly assigned to the MJDs for PSRs J1819$-$1458 and J1754$-$30 respectively, suggesting that the significance of the peak may be underestimated by the Lomb--Scargle algorithm. Given eight trials, if all of the RRATs had no periodicities, we would expect to find one periodicity with significance greater than 88\\%. However, two RRATs with significances greater than 95\\% are not expected. The predicted timescales for refractive interstellar scintillation of 117 days for PSR J1819$-$1458  and 21 days for PSR J1754$-$30 are much smaller than the reported periodicities in pulse detection rates, indicating that these periodicities are not likely due to scintillation. However, these periodicities are roughly half of the total data span.  Therefore, a longer time span of observations is necessary to confirm them as significant.  There is evidence for changes in period derivative associated with changes in pulse detection rate for PSR J1819$-$1458 \\citep{lmk09}. The peak in the rate immediately following the first (and largest) glitch at MJD 53926 \\citep{lmk09} hints at a correlation between glitches and emission properties. If confirmed in future glitches, this correlation might suggest a link with the magnetars, for which radiative changes often accompany glitches \\citep[e.g.][]{dkg+08} or the class of mode changing pulsars whose period derivative undergoes quasi--periodic changes \\citep{lhk+10}.  All of the sources exhibit random pulse distributions on single days. This is similar to that observed for giant pulses \\citep{kbm+06}. However, the pulses from the RRATs are wider than those of observed giant pulses and the pulse amplitude distributions are different \\citep{mlk+09,mmk+10}. This therefore supports the idea of Weltrevrede et al. (2006), who suggest that the RRATs could be normal, but more distant, pulsars like B0656+14 which emits very bright, narrow pulses and a distribution of weaker, broader pulses.   Over longer time spans, the pulse sequences  of six of the RRATs show evidence for non--random behaviour. These same RRATs show significant periodicities in arrival times. The external influence from, e.g., an asteroid belt \\citep{cs08} could explain this through uneven distributions of material. However, it is difficult to relate this long term non--randomness to either normal or nulling pulsars as, to our knowledge, this has not been explored for either of these source classes. We do not find any obvious relationship between the randomness and derived spin--down properties such as age, period, and surface dipole magnetic field.   Nulling pulsars tend to have on--off states which persist over more than one pulse period \\citep{wmj+06} whereas the large majority of the RRAT pulses occur singly, aside from a few pulses from PSRs J1819$-$1458 and J1317$-$5759. Redman \\& Rankin (2009) showed that the majority of pulsars null non--randomly and several studies (e.g. Rankin \\& Wright 2008 and Herfindel \\& Rankin 2009) have revealed periodicities in the null cycles for some pulsars. The RRATs have quite different properties in these respects, but because the nulls of some pulsars (e.g. B0834$+$06, B1612$+$07, and B2315$+$21; \\citet{rr09}) are random and single--pulse nulls are occasionally observed, we cannot exclude the possibility of the RRATs being an extreme case of nulling. This is supported by the similar ages and spin periods of many RRATs and nulling pulsars and the strong correlation between nulling and pulsar age \\citep{wmj+06}.  In summary, we detect highly significant periodicities in the arrival times of six of the RRATs. We detect no highly significant periodicities in the long--term pulse detection rates for the RRATs, although we find tentative periodicities for PSRs J1819$-$1458 and J1754$-$30. All of the RRATs show random behavior on a single day and most of the RRATs show non--random behavior on long timescales. Most of the RRATs emit pulses singly, but a few do show evidence for clustering of pulses. It is clear that there are periodicities in the pulse arrival times for these objects. The cause of these could be circumstellar material, non--radial oscillations, or another process. Radio monitoring over longer time spans and observations at other wavelengths may be useful to further understand the reasons for the RRATs' unusual emission. More theoretical work and further studies of the randomness of pulse emission in normal pulsars with different properties is also necessary."
    },
    "1107/1107.5116_arXiv.txt": {
        "abstract": "{We present a deep optical spectrum of TN J0924$-$2201, the most distant radio galaxy at $z = 5.19$, obtained with FOCAS on the Subaru Telescope. We successfully detect, for the first time, the \\ion{C}{iv}$\\lambda$1549 emission line from the narrow-line region (NLR). In addition to the emission-line fluxes of Ly$\\alpha$ and \\ion{C}{iv}, we set upper limits on the \\ion{N}{v} and \\ion{He}{ii} emissions. We use these line detections and upper limits to constrain the chemical properties of TN J0924$-$2201. By comparing the observed emission-line flux ratios with photoionization models, we infer that the carbon-to-oxygen relative abundance is already [C/O] $> -0.5$ at a cosmic age of $\\sim 1.1$ Gyr. This lower limit on [C/O] is higher than the ratio expected at the earliest phases of the galaxy chemical evolution, indicating that TN J0924$-$2201 has already experienced significant chemical evolution at $z = 5.19$.} ",
        "introduction": "%  The chemical evolution of galaxies is closely related to their past star-formation history. Therefore the investigation of the chemical evolution of galaxies is crucial for understanding the formation history of galaxies. Observationally, the chemical evolution of galaxies has been studied by measuring the metallicity of galaxies at various redshifts. The gas-phase metallicity of star-forming galaxies in the local universe can be measured fairly easily by means of optical emission-line spectra (e.g., [\\ion{O}{ii}]$\\lambda$3727, H$\\beta$, [\\ion{O}{iii}]$\\lambda$5007, H$\\alpha$, and [\\ion{N}{ii}]$\\lambda$6584; see Nagao et al. 2006c and references therein). However, in the early universe at $z > 1$, these diagnostics are more difficult to use because these emission lines are often very faint at $z > 1$ and are redshifted into the near-infrared, where sensitive spectroscopic observations are more challenging.  Active galactic nuclei (AGNs), thanks to their high luminosity, offer an interesting alternative to investigate the metallicity in high-$z$ galaxies. The AGN are characterized by a wealth of bright emission lines in the wavelength range from rest-frame ultraviolet to the infrared, arising from gas clouds photoionized by their central engines. In particular, the ultraviolet lines, e.g., \\ion{N}{v}$\\lambda$1240, \\ion{C}{iv}$\\lambda$1549, \\ion{He}{ii}$\\lambda$1640, and \\ion{C}{iii}]$\\lambda$1909, are useful tools to measure the metallicity in high-$z$ AGNs, because we can observe these lines in the optical. The metallicity of AGNs has been extensively studied (see Hamann \\& Ferland 1999 for a review), in particular by focusing on broad-line region (BLRs). Nagao et al. (2006a) measured emission-line flux ratios of SDSS quasars, e.g., \\ion{N}{v}/\\ion{C}{iv}, (\\ion{Si}{iv}+\\ion{O}{iv}])/\\ion{C}{iv}, \\ion{Al}{iii}/\\ion{C}{iv}, and \\ion{N}{v}/\\ion{He}{ii}, which are sensitive to metallicity. The authors found that there is no redshift evolution of the BLR metallicity in the redshift range of $2.0 < z < 4.5$. Using near-infrared spectroscopic observations of six luminous quasars at $5.8 < z < 6.3$, Jiang et al. (2007) found no strong evolution of the BLR metallicity up to $z \\sim 6$. Recently, Juarez et al. (2009) investigated the BLR metallicity of a sample of 30 quasars in the redshift range of $4.0 < z < 6.4$ through the emission-line flux ratio (\\ion{Si}{iv}+\\ion{O}{iv}])/\\ion{C}{iv} and found that the metallicity is very high even in quasars at $z \\sim 6$. These results indicate that the major epoch of chemical evolution in AGNs is at $z > 6$. However, AGN broad lines originate in a very small region in galactic nuclei ($R_{\\rm BLR} < 1$ pc; e.g., Suganuma et al. 2006), which may have evolved more rapidly and may be not representative of the metallicity in their host galaxies. Moreover, most studies have reported that the BLR metallicity is significantly higher than the solar value ($Z_{\\rm BLR} > Z_\\odot$; e.g., Hamann \\& Ferland 1992; Dietrich et al. 2003; Nagao et al. 2006a; Jiang et al. 2007), reaching as much as $Z_{\\rm BLR} \\sim 15 Z_\\odot$ in the most extreme cases (Baldwin et al. 2003b; Bentz et al. 2004). These extremely high metallicities are very hard to reconcile with galaxy chemical evolutionary models (e.g., Hamann \\& Ferland 1993; Ballero et al. 2008) and therefore probably reflect the peculiar evolution of the nuclear region.  To investigate the metallicity on AGN galactic scales, some studies have focussed on the narrow line region (NLR). In contrast to the BLRs, the typical size of the NLRs is comparable to the size of the host galaxies ($R_{\\rm NLR} \\sim 10^{1-4}$ pc; e.g., Bennert et al. 2006a, 2006b). The mass of the NLRs is about $M_{\\rm NLR} \\sim 10^{4-8} M_\\odot$, which is much higher than that of BLRs ($M_{\\rm BLR} \\sim 10^{2-4} M_\\odot$; see, e.g., Baldwin et al. 2003a). However, the metallicity of the NLR in high-$z$ AGN can be constrained only in type-2 AGNs, whose BLR and strong continuum emission are obscured and therefore allow the detection and careful measurement of the narrow lines. Because there are only a few optically-selected type-2 quasars discovered at $z > 1$, many studies of the NLR metallicity at $z > 1$ have focussed on high-$z$ radio galaxies (HzRGs; e.g., De Breuck et al. 2000; Iwamuro et al. 2003; Nagao et al. 2006b; Humphrey et al. 2008; Matsuoka et al. 2009). By studying the emission-line flux ratios of \\ion{N}{v}/\\ion{C}{iv} and \\ion{N}{v}/\\ion{He}{ii}, De Breuck et al. (2000) found that the typical NLR metallicity of HzRGs is roughly $0.4 Z_\\odot < Z_{\\rm NLR} < 3.0 Z_\\odot$. On the contrary, Nagao et al. (2006b) reported that HzRGs do not show any redshift evolution of the NLR metallicity in the redshift range $1.2 < z < 3.8$ when exploiting a metallicity diagnostic diagram involving \\ion{C}{iv}, \\ion{He}{ii}, and \\ion{C}{iii}], which allows one to separate the metallicity and ionization parameter dependencies. Matsuoka et al. (2009) confirmed the absence of any significant metallicity evolution of the NLRs in HzRGs up to $z \\sim 4$ by using a sample of 57 HzRGs at $1 < z < 4$. This result suggests that the main epoch of major evolution of the NLR metallicity occurred at even higher redshifts; it is therefore crucial to investigate the metallicity evolution in HzRGs at $z > 4$.  In this paper we focus on the most distant radio galaxy, TN J0924$-$2201 at $z = 5.19$, discovered by Van Breugel et al. (1999). TN J0924$-$2201 is currently the only target that allows us to investigate the NLR metallicity at $z > 5$. In previous studies, Venemans et al. (2004) found that TN J0924$-$2201 is located in a protocluster environment: the density of Ly$\\alpha$ emitters around this object is comparable to the density in protoclusters found around HzRGs at $z < 4$ (see also, Overzier et al. 2006). By analyzing the Spitzer data, Seymour et al. (2007) derived the stellar mass of TN J0924$-$2201, i.e., 10$^{11.10} M_\\odot$, comparable with the stellar masses of HzRGs at $1 < z < 4$. Klamer et al. (2005) reported that TN J0924$-$2201 may be a young forming galaxy based on its low dust-to-gas ratio (inferred from its $L_{\\rm FIR}/L_{\\rm CO}$ ratio). This suggests that TN J0924$-$2201 might be undergoing a significant chemical evolution, which can be verified by studying the metallicity of its NLR. However, since previous optical spectroscopic observations of this radio galaxy have identified only the Ly$\\alpha$ emission line (i.e., Van Breugel et al. 1999; Venemans et al. 2004), we have obtained a very deep optical spectrum of TN J0924$-$2201 with FOCAS at the Subaru Telescope aimed at detecting additional narrow lines. In the rest of the paper, we adopt a concordance cosmology with ($\\Omega_{\\rm M}$, $\\Omega_\\Lambda$) = (0.27, 0.73) and $H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "%  To investigate the chemical evolution of HzRGs at $z > 5$, we investigated the rest-frame UV spectrum of TN J0924$-$2201 and obtained the following main results. \\begin{enumerate} \\item We detected \\ion{C}{iv} emission in TN J0924$-$2201 at $z = 5.19$; this is the first detection of metal emission lines from NLR gas clouds at $z > 5$, suggesting that a significant amount of carbon exists even at $z > 5$ in the early universe. \\item The measured Ly$\\alpha$/\\ion{C}{iv} flux ratio is lower than that of HzRGs at $3 < z < 5$; this can possibly be ascribed to weaker star-formation activity and/or heavier IGM absorption on Ly$\\alpha$. \\item By comparing observational limits on emission-line flux ratios (and in particular \\ion{N}{v}/\\ion{C}{iv} and \\ion{C}{iv}/\\ion{He}{ii}) with photoionization models, we inferred that the carbon is already relatively high in TN J0924$-$2201, more specifically [C/O] $> -0.5$, suggesting that TN J0924$-$2201 is not in an extremely young phase, but likely already a few hundred million yrs old. \\end{enumerate}"
    },
    "1107/1107.0727_arXiv.txt": {
        "abstract": "Model-independent parametrisations of modified gravity have attracted a lot of attention over the past few years and numerous combinations of experiments and observables have been suggested to constrain the parameters used in these models. Galaxy clusters have been mentioned, but not looked at as extensively in the literature as some other probes. Here we look at adding galaxy clusters into the mix of observables and examine how they could improve the constraints on the modified gravity parameters. In particular, we forecast the constraints from combining Planck satellite Cosmic Microwave Background (CMB) measurements and Sunyaev--Zeldovich (SZ) cluster catalogue with a DES-like Weak Lensing (WL) survey. We find that cluster counts significantly improve the constraints over those derived using CMB and WL. We then look at surveys further into the future, to see how much better it may be feasible to make the constraints. ",
        "introduction": "Einstein's General Relativity (GR) is one of the principal ingredients of modern cosmology. Indeed, it could be argued that it was only with the development of GR that cosmology really became a part of physics. Nonetheless, it is our job as physicists to continue to test even the most fundamental pillars of cosmology in order to refine, improve and further justify our model of the universe. There are also fundamental reasons for considering different theories of gravity: GR is inconsistent with quantum mechanics and the search for a theory of `Quantum Gravity' is one of the holy grails of modern physics.  It has proved to be difficult to test GR outside of the solar system, particularly as the effects of a different theory of gravity could be degenerate with behaviour induced by different constituents of the universe. This is the case with current observations that suggest the presence of some form of dark matter and dark energy. Dark matter is required to explain galaxy rotation curves, galaxy lensing, nucleosynthesis, acoustic oscillations in the CMB, and the growth of structure in the Universe to name a few. The requirement for dark energy is underpinned by observations of the background expansion rate and large scale structure measurement.  The question has often been raised as to whether these effects could be due to a modified gravity theory. Proposals for such a theory include a number of $f(R)$ theories, the Dvali Gabadadze Porrati (DGP) model \\cite{dgp}, conformal gravity \\cite{cfg}, Modified Newtonian dynamics (MOND) \\cite{mond} and its covariant, relativistic extensions \\cite{teves}, and Einstein Aether theories \\cite{eatheories}, (see also the recent review \\cite{mgreview} for an exhaustive list of candidates). It is possible, if GR is indeed the correct theory of gravity, that this debate will only be settled by the non-gravitational detection of the dark matter and/or dark energy. However, since we do not know whether or not GR is indeed the correct theory of gravity, it seems reasonable to consider which observations could allow us to detect modified gravity.  More importantly, are there ways to test deviations from GR in a model independent way? There are several advantages to a model independent approach; some alternatives to GR do exist but there is no complete theory of, for example, quantum gravity to draw on. Also, there are no `stand-out' candidates that are universally considered to be strong alternatives. Another advantage of a model independent approach is that the results do not rely on model selection techniques such as $\\chi^2$ per degree of freedom or other ways of choosing between competing theories. Thus, they also do not rely on us having {\\sl the} correct theory to hand. A result that is inconsistent with the GR based, concordance cosmology will be unambiguous and therefore a strong motivator to develop alternative theories, as well as possibly giving us a clue as to the nature of these theories.  There have been many recent studies looking at model independent tests of the dark energy/cold dark matter ($\\Lambda$CDM) paradigm in GR. These can be roughly split into two categories: consistency checks of the $\\Lambda$CDM assumption \\cite{song09,shapiro10,acqua10} and those that introduce new parameters to evaluate the level of deviation from GR (see \\cite{mgreview} for a complete review of the literature). These parametrisations and consistency checks have considered the majority of cosmological observations: weak lensing, the CMB, particularly through its Integrated Sachs Wolfe (ISW) effect, Baryon Acoustic Oscillations (BAO), SN1a Supernovae luminosity distance observations, cluster counts, and galaxy redshift and peculiar velocity surveys. The purpose is normally to constrain both the background expansion history and the perturbations, or growth of structure, around the background, with some observables being sensitive to both. The constraints are then combined. This approach works well because the growth and expansion are determined by the same quantities under the concordance cosmology and this is the basis for the consistency checks. A parameter that can be used without modifying gravity is $\\gamma$, a parametrisation of the growth index \\cite{linder05,wang98}: \\begin{equation} \\frac{d\\ln D(k,a)}{d\\ln a}=\\Omega_m(a)^{\\gamma}\\,. \\end{equation} Here, $D(k,z)=\\delta(k,z)/\\delta(k,z=\\infty)$ where $\\delta$ is the matter density contrast. For $\\Lambda$CDM, $\\gamma=0.55$ is a good fit to the growth. With the modified gravity parametrisations, the parameters often relate to the two gravitational potentials $\\Psi$ and $\\Phi$ appearing in the perturbation to the Friedmann Roberston Walker (FRW) metric. Since different observations depend on different combinations of these potentials, combining several experiments gives the best constraints.  Constraints from current data have been examined (see \\cite{mgreview} for a review) and the general conclusion appears to be that the concordance cosmology is consistent with all of the current data. However, the data available today does not have enough constraining power to rule out even relatively significant modifications to gravity and they are certainly not precise enough to distinguish {\\sl between} modified gravity theories. Work has also gone into forecasting future constraints on a number of theories and/or parametrisations of the modifications to GR \\cite{gbz10b,gbz09,dossett10,daniel10b,serra09,guzik10,kosow09,heavens07,acqua10}. The consensus is that future surveys will greatly improve prospects with the expectation that a number of theories competing with GR will be ruled out.  Now is a good time to consider these issues as we are in an era where the observational front is rapidly advancing in the field. Most, if not all, of the current and future measurements can be brought to bear on the issue of modified gravity. We have had pioneering ground and space-based CMB experiments over the last 20 years and now await the data from the Planck satellite \\cite{planck}. Weak lensing surveys of cosmic shear have also matured into a precise observational tool \\cite{cfhtls} and future, planned surveys promise to bring these measurements to the fore front of the data landscape with large scale surveys underway or in the development stage (e.g.  PANSTARRS \\cite{panstarrs}, DES \\cite{des}, and LSST \\cite{lsst}). In addition large scale surveys of galaxy redshifts have already been carried out (SDSS \\cite{sdss}, and 2DFGRS \\cite{2dfgrs}) with even larger and deeper ones targeting BAO measurements to come.  So far, we have only mentioned modified gravity as an alternative to dark energy. The other alternative under consideration is whether the assumptions of homogeneity and isotropy are justified and hence whether using the FRW metric itself is justified. It is also important to note that measurements of the background expansion alone cannot distinguish between dark energy or modified gravity \\cite{bert08} and it has also been argued  that measurements of perturbations may also suffer from this degeneracy if sufficiently complex models of perturbed dark energy are allowed \\cite{kunz06}.  In this work, we will only consider modified gravity and will look at how well some future experiments, particularly DES and CMB measurements and SZ cluster counts from Planck and its successors, will combine to constrain certain modified gravity parameters that characterise potential deviations from GR.  In this {\\sl paper} we investigate how the combination of future observations of CMB, weak lensing (WL), and cluster counts (CC) are able to constrain the model independent parametrisation of modified gravity theories. We restrict ourselves to the simplest form of modified gravity with at most a linear redshift dependence in the modification and no scale dependence. These assumptions are fairly restrictive in terms of the physical mechanism that could underpin the modified phenomenology but provide a simple starting point for investigation into future constraints from different observables. In Section~\\ref{sec:gravity} we briefly review the model independent parametrisation of modified gravity used in this work. In Section~\\ref{sec:obs} we describe the three observables used in our forecasts. In Section~\\ref{sec:forecasts} we review the Fisher matrix formalism used in our calculations and the experimental `Stages' considered in our forecasts. Our results are presented in Section~\\ref{sec:results} and we conclude with a discussion of the results in Section~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  Future generations of large area surveys hold much promise in testing the validity of GR. In this work we have examined the impact, on MGP constraints, of future cosmic shear and cluster count data used in combination with CMB measurement. We have found that this combination can constrain the MGPs to sub--percent accuracy. In particular the inclusion of cluster counts, which are highly sensitive to any change in the growth of the matter perturbations, adds a strong refinement in the search for any deviations from the standard GR values of the MGPs. However we have also shown that the inclusion of a simple model for calibration uncertainties in the counts can affect the errors.  Our Fisher matrix analysis has shown that the MGPs may be {\\sl statistically} constrained to fractional levels of below a percent. This clearly supercedes the current modeling accuracy of weak lensing and cluster abundances in particular. In the case of weak lensing, even at large scales, non-linear contributions to the matter power spectrum are not modeled accurately enough whilst in the cluster case the mass function formalism is certainly not accurate at the 1\\% level. Any analysis of actual data will necessarily have to include a marginalisation over these and other modeling uncertainties. This can be done either by including the model uncertainties in the final Fisher matrix and marginalising by cutting out the directions in its inverse or by including the model parameters as extra directions in MCMC methods. Future MGP constraints may well be limited by such systematic effects rather than the statistical limits of the observations.  There are, of course, other observables that can constrain the growth of matter perturbations such as large scale structure, peculiar velocities, redshift space distortions and future high redshift 21cm surveys. These could all replace our choice of cluster counts as a signal used to break the degeneracy between gravitational potentials. All of these, however, present a number of problems either in their interpretation as biased tracers or in the technological challenges involved in large scale surveys. The observations of clusters on the other hand is a relatively simple procedure and SZ surveys of clusters are a necessary by--product of current and future CMB experiments. The only difficulty involved in this signal is the theoretical modelling required to use them to obtain constraints due to the uncertainty in the choice of mass functions to be used for theories of modified gravity.  In obtaining our forecasted constraints we have only employed the simplest description of phenomenologically modified gravity. The model building for such theories has already shown that there there is much more complexity to explore. All models must incorporate a screening mechanism to satisfy standard solar system and laboratory tests of gravity. Screening mechanism are only one way in which modified gravity models may include significant scale dependence, which we have not taken into account in this work and may provide more testable predictions and therefore more ability to differentiate between models. Concrete models may also lead to modifications to the background expansion, which we have not explored here. This can also lead to additional constraints on model parameters.  Allowing for a very general scale dependence of the MGPs will inevitably introduce a scale dependence into the constraints. In particular our combination of data generates constraints based on large scales, due to the ISW effect, and a fairly narrow range of scales of around 10 Mpc from cluster counts. If there is a significant scale dependence of the MGPs between the two scales then the combination of data considered here would give stronger constraints on some scales than others. A complete treatment would include a Fisher eigenmode analysis for an extended range of parameters describing the scale dependence of the MGPs \\cite{gbz10,gbz10b,gbz09}. We leave this for future work for the combination of data considered here.  The experimental outlook is promising for the observations we have dealt with in this work. Over the next 5-10 years, deviations from GR should be well constrained, and the concordance cosmology will either be more secure or may even have undergone a paradigm shift. A null result would support the concordance cosmology, a conclusion that would be even stronger if dark matter had been detected non-gravitationally by then. A detection of a deviation from GR would be potentially more interesting but would require a completely new theoretical framework and trigger a search for an underlying model for the modifications.  Of course, a dark energy model with perturbations may turn out to fit the data just as well and it is not clear at this point whether this degeneracy will ever be broken by observations.  {\\sl Acknowledgements} We thank Filipe Abdalla for pointing out the work on self calibration of clusters to us. This work was supported by an STFC studentship.  \\appendix"
    },
    "1107/1107.2208_arXiv.txt": {
        "abstract": "Modelling of adiabatic gravity wave propagation in the solar atmosphere showed that mode conversion to field guided acoustic waves or Alfv\\'en waves was possible in the presence of highly inclined magnetic fields. This work aims to extend the previous adiabatic study, exploring the consequences of radiative damping on the propagation and mode conversion of gravity waves in the solar atmosphere.   We model gravity waves in a VAL-C atmosphere, subject to a uniform, and arbitrarily orientated magnetic field, using the Newton cooling approximation for radiatively damped propagation.  The results indicate that the mode conversion pathways identified in the adiabatic study are maintained in the presence of damping.  The wave energy fluxes are highly sensitive to the form of the height dependence of the radiative damping time.  While simulations starting from 0.2 Mm result in modest flux attenuation compared to the adiabatic results, short damping times expected in the low photosphere effectively suppress gravity waves in simulations starting at the base of the photosphere.   It is difficult to reconcile our results and observations of propagating gravity waves with significant energy flux at photospheric heights unless they are generated \\emph{in situ}, and even then, why they are observed to be propagating as low as 70 km where gravity waves should be radiatively overdamped. ",
        "introduction": "\\label{intro}    Recent multi-height observations of low frequency oscillations in the low solar chromosphere suggest that the energy flux carried by upward propagating gravity waves, with frequencies between 0.7 and 2.1 mHz, comfortably exceeds the co-spatial acoustic wave flux \\citep{straus08}.  In our previous paper  \\citep{newington10}, (referred to henceforth as paper I), we explored the propagation, reflection and mode conversion of gravity waves in a VAL C  related atmosphere, permeated by uniform, inclined magnetic field.  The significant finding of that study was that in regions of highly inclined magnetic field, gravity waves experience mode conversion to up-going (field-guided) acoustic or Alfv\\'en waves.  While acoustic waves are likely to shock before reaching the upper chromosphere, Alfv\\'en waves can propagate to greater atmospheric heights, perhaps contributing to the observed coronal Alfv\\'enic oscillations \\citep{pontieu07,tomczyk07}.  A limitation of the work presented in paper I was the assumption of adiabatic wave propagation, which is known to be invalid in the photosphere and low chromosphere.   A more realistic investigation of the propagation and mode conversion of gravity waves in the solar atmosphere would include the effects of radiative damping.  Radiatively damped atmospheric gravity waves have been considered by a small number of authors previously, (see \\citealt{souffrin66}, \\citealt{stix70}, and the exhaustive study by \\citealt{mihalas82}), but all of those studies have been purely hydrodynamic in nature, with no applied magnetic field, hence there is no possibility of mode conversion.  Our principal objective in this work is to explore the consequences of radiative damping on the mode conversion of gravity waves, and this requires a magnetohydrodynamic (MHD) treatment.  In this paper we extend our earlier MHD study, relaxing the assumption of adiabatic wave propagation by incorporating radiative damping using the Newton cooling approximation.  We recognise that this approximation is strictly valid only for optically thin perturbations of a homogeneous, infinite,  and isothermal gas, and that none of these conditions are realised in the photosphere and chromosphere.  The advantage of using this simple treatment is that the same mathematical tools used in paper I can be employed with minor modifications.  The insight gained from this simple model may direct our attention to cases that are interesting enough to warrant a more thorough treatment of the effects of radiation.  The questions central to our investigation are: \\textit{How are the mode conversion pathways affected by radiative damping? Do we still get appreciable mode conversion to Alfv\\'en waves when damping is present?}  This paper is organised as follows:  Section 2 presents the atmospheric model and the mathematical tools used in this investigation, namely the damped dispersion relation and numerical solution of the damped linear MHD wave equations.  Section 3  presents the results.  Dispersion diagrams reveal insights into the mode conversion pathways in $z$-$k_z$ phase space, and numerical integration of the wave equations quantify the acoustic and magnetic fluxes.  The conclusions are stated in Section 4. ",
        "conclusions": "\\label{conclusions} The main conclusions we draw from the results of our investigation are:  \\begin{enumerate} \\item The primary effect of radiative damping in our model is a decrease in the measured wave energy fluxes obtained, with respect to the adiabatic results.  The extent of the flux attenuation is highly sensitive to the radiative damping times.  \\item Damping does not have a great effect on the mode conversion pathways in the $z-k_z$ plane. Specifically, direct coupling to Alfv\\'enic disturbances higher up in the chromosphere seems to survive (in diminished form) the introduction of radiative losses.  \\item Short damping times very effectively extinguish gravity waves.  This  is most likely to be a problem in the low photosphere, where radiative damping times are expected to be far shorter than wave periods, rendering the waves overdamped.  Reconciling observations of propagating photospheric gravity waves \\citep{straus08,stodilka08} with this seems an impossibility unless they are being generated there \\emph{in situ}.\\footnote{In Earth and planetary atmospheres, it is known that primary gravity waves can produce secondary gravity waves by nonlinear processes such as wave breaking and the production of turbulent cascades \\citep{lane06}, and non-linear wave-wave interactions \\citep{dong88, huang07}. Such processes do not seem to be available to us in the low solar atmosphere.}   \\cite{komm91} suggest that gravity waves are generated by decay of granular motions in the mid photosphere. \\cite{stodilka08} found that overshooting convection starts at around $100$ km and apparently detected gravity waves above this height.  Gravity wave excitation by penetrative convection (albeit in the context of the convection zone/radiative core interface) has been modelled numerically \\citep{din03,din05} and found to be effective.  \\item Even accepting \\emph{in situ} gravity wave generation, it is not clear that \\cite{straus08} see \\emph{propagating} gravity waves as low as 70 km, where they should be substantially overdamped, notwithstanding the rather arbitrary form of our radiative decay time $\\tau(z)$. Indeed, their own numerical simulations using the (nonmagnetic) radiation hydrodynamics \\textsf{CO$^\\mathsf{5}$BOLD} code, which has a much more sophisticated treatment of radiation than we have instituted, is dominated by convection at $\\sim70$ km, with a negligible gravity wave flux there (Straus, private communication). Gravity waves are progressively generated as height increases, reaching their full strength at around 300 km, with the gravity wave being `free' only above this. This justifies our decision to model gravity waves in $z>200$ km; the region below 200--300 km is where the waves are driven.   \\end{enumerate}"
    },
    "1107/1107.0661_arXiv.txt": {
        "abstract": "{ We present medium-resolution ($\\lambda / \\Delta\\lambda \\sim 20,000$) ultraviolet spectra covering the 1155--1760 \\AA\\ spectral range of the Seyfert 1 galaxy \\mrk509 obtained using the Cosmic Origins Spectrograph (COS) on the Hubble Space Telescope (HST). Our observations were obtained simultaneously with a Low Energy Transmission Grating Spectrometer observation using the {\\it Chandra} X-ray Observatory, and they are part of a multiwavelength campaign in September through December 2009 which also included observations with {\\it XMM-Newton}, Swift, and INTEGRAL. Our spectra are the highest signal-to-noise observations to date of the intrinsic absorption components seen in numerous prior ultraviolet observations. To take advantage of the high signal-to-noise ratio, we describe special calibrations for wavelength, flat-field and line-spread function corrections that we applied to the COS data. We detect additional complexity in the absorption troughs compared to prior observations made with the Space Telescope Imaging Spectrograph (STIS) on HST. We attribute the UV absorption to a variety of sources in \\mrk509, including an outflow from the active nucleus, the interstellar medium and halo of the host galaxy, and possible infalling clouds or stripped gaseous material from a merger that are illuminated by the ionizing radiation of the active nucleus. Variability between the STIS and COS observation of the most blue-shifted component (\\#1) allows us to set an upper limit on its distance of $< 250$ pc. Similarly, variability of component 6 between FUSE observations limits its distance to  $<$ 1.5 kpc. The absorption lines in all components only partially cover the emission from the active nucleus with covering fractions that are lower than those seen in the prior STIS observations and are comparable to those seen in spectra from the Far Ultraviolet Spectroscopic Explorer (FUSE). Given the larger apertures of COS and FUSE compared to STIS, we favor scattered light from an extended region near the active nucleus as the explanation for the partial covering. As observed in prior X-ray and UV spectra, the UV absorption has velocities comparable to the X-ray absorption, but the bulk of the ultraviolet absorption is in a lower ionization state with lower total column density than the gas responsible for the X-ray absorption. We conclude that the outflow from the active nucleus is a multiphase wind. } ",
        "introduction": "Outflows from galaxies powered by an active galactic nucleus (AGN) may play an important role in the chemical enrichment of the intergalactic medium (IGM). The most distant AGN known, quasars at redshifts of 6 or more, have abundances at solar or super-solar levels \\citep{Hamann93, Hamann99, Pentericci02, Barth03, Dietrich03, Freudling03, Juarez09}. Although the chemical enrichment of the host likely predates the presence of the active nucleus \\citep{Simon10}, winds powered by the AGN could return this metal-enriched material to the IGM \\citep{Furlanetto01, Cavaliere02, Germain09, Hopkins10, Barai11}. Their presence and feedback may also have a significant impact on the evolution of their host galaxies \\citep{Silk98, Scannapieco04, Granato04, DiMatteo05, Hopkins08, Somerville08}. While some fraction of these outflows in low-luminosity AGN may not escape their host galaxy, at least as measured in the local universe \\citep{Das05, Ruiz05, Das07}, the impact of the outflow on lower-density portions of the host interstellar medium (ISM) could provide a significant transport mechanism for the enrichment of the surrounding environment \\citep{Hopkins10}. From constraints provided by the X-ray background, such lower luminosity AGN may dominate the population of active galaxies in the early universe \\citep{Treister09, Treister10}.  Nearby AGN provide local analogs that can help us to understand the mechanics, energetics, and chemical enrichment patterns that may play a significant role in cosmic evolution at high redshift. More than half of low-redshift AGN exhibit blue-shifted UV or X-ray absorption features indicative of outflowing gas \\citep{Crenshaw03, Dunn07, Cappi09, Tombesi10}. Understanding the geometry and the location of the outflow relative to the active nucleus is a key to making an accurate assessment of the total mass and the kinetic luminosity of the outflow. Distance determinations are particularly difficult. Using density-sensitive absorption lines to establish the gas density, in combination with photoionization models that reproduce the observed relative column densities can provide precise distance estimates. These measures have ranged from tens of parsecs in NGC 3783 \\citep{Gabel05b} and NGC~4151 (\\citealt{Kraemer06}, if one uses the correct critical density for \\ciii*), and up to kiloparsec scales in some quasars and AGN \\citep{Hamann01, Scott04a, Edmonds11}. UV and X-ray observations of Mrk~279 \\citep{Scott04a, Arav07, Costantini07} measured absolute abundances in the outflow of a local AGN for the first time, and they show evidence for enhanced CNO abundances that suggest contributions from massive stars and AGB stellar winds.  To improve upon these prior studies, we have conducted a multiwavelength campaign of coordinated X-ray, UV, and optical observations of the nearby luminous Seyfert 1 galaxy \\mrk509 (z=0.034397; \\citet{Huchra93}). A complete overview of the campaign is given by \\citet{Kaastra11a}. \\mrk509 is an ideal object for study due to its high flux, moderate luminosity that rivals that of QSOs \\citep{Kopylov74}, and deep, well structured absorption troughs. The blue-shifted absorption indicative of a nuclear outflow has been known since the earliest UV spectral observations \\citep{Wu80, York84}. The outflow is also apparent in the blueshifted {[O\\,{\\sc iii}]} emission detected across the face of \\mrk509 by \\citet{Phillips83}. More recently \\citet{Kriss00} and \\citet{Kraemer03} observed \\mrk509 at high spectral resolution in the UV with the Far Ultraviolet Spectroscopic Explorer (FUSE) and the Space Telescope Imaging Spectrograph (STIS) on the Hubble Space Telescope (HST), respectively. The STIS observations were simultaneous with earlier {\\it Chandra} X-ray grating observations \\citep{Yaqoob03} that also detect blue-shifted X-ray absorption lines. These prior UV observations show the presence of more than seven distinct kinematic components, none of which could be unambiguously associated with the absorption detected in the X-ray, although both the X-ray and UV absorption do overlap in velocity.  Our newer observations using the Cosmic Origins Spectrograph (COS) on HST achieve higher signal-to-noise ratios (S/N) and provide insight into long-term changes in the absorbing gas. Our COS spectra provide a higher-resolution view of the velocity structure in the outflowing gas, and sample lower ionization ionic species than are present in the X-ray spectra. At high spectral resolution and high S/N, we can use velocity-resolved spectral coverage of the \\nv\\ and \\civ\\ doublets to solve for the covering fractions and column densities in the outflowing gas \\citep{Arav07}. Using our COS observations of \\lya\\ together with the archival FUSE observations of the higher-order Lyman lines, we can provide an absolute reference for abundances relative to hydrogen, which cannot be measured directly in the X-ray spectra. These higher S/N spectra combined with the higher S/N X-ray spectra permit a more detailed examination of the links between the UV and the X-ray absorbing gas in the outflow.  This paper describes the COS observations that are part of our multiwavelength campaign on \\mrk509, and presents an initial interpretation of the results. In \\S2 we describe our COS observations and our data reduction methods, including enhanced calibrations for the COS data and corrections for the line-spread function. In \\S3 we describe our analysis of the UV spectra and compare them to prior STIS and FUSE observations of \\mrk509. Finally, in \\S4 we discuss the physical implications of our observations for the geometry and location of the absorbing gas in \\mrk509 and its physical conditions. We end with a summary of our conclusions in \\S5. ",
        "conclusions": "We have presented HST/COS observations of the Seyfert 1 galaxy \\mrk509 obtained as part of an extensive multiwavelength campaign \\citep{Kaastra11a} on this galaxy in late 2009. Our UV spectra were obtained simultaneously with {\\it Chandra} LETGS spectra \\citep{Ebrero11a}. The campaign also included immediately prior observations with {\\it XMM-Newton}, Swift, and INTEGRAL. Our spectra cover the 1155--1760~\\AA\\ wavelength range at a resolution of $\\sim15~\\kms$, and they are the highest signal-to-noise observations to date of the intrinsic absorption components of \\mrk509. This enables us to trace additional complexity in the absorption troughs compared to prior STIS observations \\citep{Kraemer03}. As part of our analysis, we have also examined archival FUSE observations, which provide additional information on variability characteristics of the absorbers.  In our COS spectra we identify a total of 14 kinematic components. Six of these are blends that represent subcomponents of previously known features. The most blueshifted portions of the \\lya\\ absorption trough (components 1 and 1b) show a significant decrease in column density compared to prior STIS observations. At the same time, component 1 in \\nv\\ increased in strength compared to the STIS observations. Using recombination timescale arguments, the variation in the \\nv\\ column density allows us to set a lower limit on the gas density of $>160~\\cc$, and an upper limit on the distance of $< 250$ pc for this most blue-shifted component. We also detect variability between the two FUSE observations from 1999 November \\citep{Kriss00} to 2000 September in component 6 at a positive velocity of $+120~\\kms$, which corresponds to the peak of the {[O\\,{\\sc iii}]} emission-line profile from the echelle spectrum of \\citet{Phillips83}. Based on the recombination time, we set a limit on the density of this component of $>6~\\cc$, and set an upper limit on its distance from the nucleus at $<1.5~\\rm kpc$.  The COS spectra also show greater residual light at the bottoms of the absorption troughs relative to the STIS spectrum. The fraction of residual light seen in COS is comparable to that seen in the FUSE spectra. We attribute this residual light to scattering from an extended region near the nucleus, which is consistent with the larger apertures used for the COS and FUSE observations relative to the STIS spectrum. Analysis of the cross-dispersion profile in the COS spectrum limits the spatial extent of this scattering region to $< 0.5 \\arcsec$, or $< 170$ pc. Scattering from such a region would be consistent with electron scattering from the base of an ionization cone in \\mrk509, and it would explain the polarization in the core of \\Halpha\\ \\citep{Young99} which is observed to be perpendicular to the radio jet \\citep{Singh92}.  The velocities of the two main absorption troughs roughly correspond to the kinematics of the detected X-ray absorption. The UV absorption in \\mrk509 arises from a variety of sources. The lowest-velocity absorption trough covers a velocity range of $-100~\\kms$ to $+250~\\kms$. The deepest portion of this trough in the UV is at the systemic velocity of \\mrk509, and it has the lowest ionization. This is also true for the X-ray absorption. We attribute this portion of the absorption to the interstellar medium or galactic halo of \\mrk509. The most redshifted portion of the trough has characteristics comparable to the high-velocity cloud Complex C in our own Milky Way \\citep{Thom08}. At a density of $0.01~\\cc$ (which we can only assume, not measure), it would have a size of roughly 1.3 kpc at a distance of 19 kpc from the center of \\mrk509.  The most blue-shifted absorption trough in \\mrk509 is an outflow from the active nucleus. The velocities in this trough of $-200~\\kms$ to $-450~\\kms$ correspond to the most highly ionized portion of the X-ray absorbing gas \\citep{Detmers11, Ebrero11a}, and they overlap with the extended emission from high-ionization gas that covers the face of \\mrk509 in the three-dimensional spectroscopy of \\citet{Phillips83}. The outflow velocities, embedded clumps of lower ionization gas, the presence of an extended scattering region, and the extended outflowing emission seen in {[O\\,{\\sc iii}]} \\citep{Phillips83} are all consistent with an origin in a multiphase thermal wind produced by the irradiation of an obscuring torus by the active nucleus."
    },
    "1107/1107.5044_arXiv.txt": {
        "abstract": "Using the Sloan Digital Sky Survey Data Release 7, we reexamine the Eastern Banded Structure (EBS), a stellar debris stream first discovered in Data Release 5 and more recently detected in velocity space by Schlaufman et al. The visible portion of the stream is $18\\arcdeg$ long, lying roughly in the Galactic Anticenter direction and extending from Hydra to Cancer. At an estimated distance of 9.7 kpc, the stream is $\\approx 170$ pc across on the sky.  The curvature of the stream implies a fairly eccentric box orbit that passes close to both the Galactic center and to the sun, making it dynamically distinct from the nearby Monoceros, Anticenter, and GD-1 streams. Within the stream is a relatively strong, 2\\arcdeg-wide concentration of stars with a very similar color-magnitude distribution that we designate Hydra I. Given its prominance within the stream and its unusual morphology, we suggest that Hydra I is the last vestige of the EBS's progenitor, possibly already unbound or in the final throes of tidal dissolution. Though both Hydra I and the EBS have a relatively high velocity dispersion, given the comparitively narrow width of the stream and the high frequency of encounters with the bulge and massive constituents of the disk that such an eccentric orbit would entail, we suggest that the progenitor was likely a globular cluster, and that both it and the stream have undergone significant heating over time. ",
        "introduction": "At least 14 stellar debris streams in the Galactic halo have now been identified in photometric surveys (see \\citet{grill2010} for a review). A similar number of dynamically cold substructures have been detected in velocity space \\citep{helmi1999, smith2009,schlaufman2009, williams2011}.  Each of these streams is interesting as a partial record of the accretion history of our Galaxy. However, and perhaps more importantly, these streams can also serve as very sensitive probes of the Galactic potential \\citep{law2009, koposov2010}. Globular cluster streams are particularly important in this respect as they are dynamically very cold \\citep{combes1999, odenkirchen2009, willett2009}. A large sample of streams will eventually enable us to constrain the distribution of dark matter in the halo in a detailed and self-consistent manner. Enlarging the sample of known streams will also increase the probability that we may detect unmistakable signs of perturbations by dark matter sub-halos \\citep{mura99, carlberg2009, yoon2010}.  In this paper we reexamine the EBS first detected by \\citet{grill2006b} using the more complete coverage available in the SDSS Data Release 7 (DR7, \\citet{abazajian2009}). We briefly describe our analysis in Section \\ref{analysis}. We characterize the EBS and and a possible progenitor in Section \\ref{discussion} and we put preliminary constraints on the orbit in Section \\ref{orbit}. We make concluding remarks Section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  The EBS stream adds to the growing list of halo streams that can be mapped over a sufficient extent that, with suitable follow-up observations, they could be used as probes of the Galactic potential. A preliminary orbit estimate shows that the EBS is unrelated to either the Anticenter or Monoceros streams.  The somewhat intermediate breadth of the stream together with its relatively high velocity dispersion suggests the possibility that the progenitor could have been more massive than the globular clusters thought to be responsible for the half dozen very cold streams discovered in the SDSS footprint to date. However, if the progenitor had been a dark matter dominated dwarf galaxy, it would be difficult to understand how it could have held onto it's dark matter envelope for any length of time in such a confined and eccentric orbit.  On the other hand, this very orbit may have subjected both the progenitor and the stream to significant heating through encounters with massive structures in the disk.   If Hydra I is indeed the progenitor of the EBS, then it is only the second probably unbound progenitor to be associated with a tidal stream. A more detailed examination of the structure and stellar kinematics in this remnant may shed new light on the end stage of tidal disruption.  Though contamination by field stars is high, Hydra I may be particularly attractive in this respect as it is four times closer to us than Bootes III \\citep{grill2009}.  Refinement of the orbit will require radial velocity and proper motion measurements of carefully selected stars along the length of the stream. Given the very low surface density of stream stars and very high field star contamination, this will necessarily be an ongoing task. In this respect, the EBS may be particularly well situated for follow-up by the upcoming spectroscopic LAMOST survey. Gaia and LSST proper motion measurements may also help us to refine the orbit and perhaps trace the stream over a much longer arc."
    },
    "1107/1107.3454_arXiv.txt": {
        "abstract": "We argue that transient systems with luminosity between novae and supernovae (SNe) are powered by gravitational energy of mass accreted onto, or a companion merges with, a main-sequence star. These transient events are termed Intermediate-Luminosity Optical Transients (ILOTs; other terms in use are Intermediate-Luminosity Red Transients and Red Novae). We show that despite the wide range of $10^{45}$--$10^{50} \\erg$, the typical energy released by ILOTs can be expressed as a function of fundamental variables $F_I(\\hbar, c, G, m_e, m_N, kT_i)$, the planck constant, speed of light, gravitational constant, electron mass, neutron mass, and ignition temperature of hydrogen. This expression explains why ILOTs are located between SNe and novae with respect to their total energy. We also put an upper limit on the power (luminosity) of ILOTs, which explains their lower luminosity than SNe. ",
        "introduction": "\\label{sec:intro}  The total energy released in supernova (SN) explosions are about six orders of magnitude higher than the total energy released in novae explosions. By \\emph{total energy} we refer here to the sum of kinetic energy of the ejected material and the total radiated energy. By \\emph{available energy} we refer to the total energy that is released in the powering process. Observations slowly filling the gap between novae and SNe (e.g., Barbary et al. 2009; Berger et al. 2009, 2011; Bond et al. 2009; Kulkarni et al. 2007a,b; Kulkarni \\& Kasliwal 2009; Ofek et. al. 2008; Rau et al. 2007; Kasliwal et al. 2011; Monard 2008; Prieto et al. 2008; Nakano 2008; Mason et al. 2010 ; Mould et al. 1990; Pastorello et al. 2010; Smith et al. 2011; Wesson et al. 2009). In our presentations at the \\emph{Intermediate-Luminosity Red Transients} meeting we argued that the total energy of these transients hints on the basic energy source (`engine') powering their eruptions. We term these outbursts Intermediate-Luminosity Optical Transients (ILOTs; other terms in use are Intermediate-Luminosity Red Transients and Red Novae).  In this summary of our presentations we show that in the same way as the total energy of SNe and novae can be explained from fundamental quantities, so is the case for ILOTs. For that we summarize in short the expressions for SNe and novae energies, and then turn to discuss ILOTs. We propose that ILOTs are powered by accreting mass from a companion onto a main-sequence (MS) star. The companion can either survive the mass transfer process, or merge with the primary in a process termed mergerburst. ",
        "conclusions": "\\label{sec:summary}  We here take the view that the systems populating the peak-luminosity region, or total energy region, between novae and SNe are not a collection of different kinds of objects. Rather, we think that the typical total energy (kinetic + radiated) of these events (termed ILOTs) points to one basic energy source: The gravitational energy of mass accreted onto, or merged with, a main-sequence star. The energy scale of ILOTs comes from the basic properties of main-sequence stars, namely, that they burn hydrogen. First, the potential on their surface is related to the hydrogen ignition temperature by equation (\\ref{eq-MS1}). Second, the available energy for a major event is the mass a companion can supply times the potential. The typical mass in a merger process is that of another main-sequence star. In a case of a mass transfer from the envelope of an evolved star a similar mass is available. In many cases less mass is transferred, and the total energy is lower even (section \\ref{sec:time}).  The power of the basic engine that leads to ILOTs is limited by viscosity that limits the duration of the process from below. This in turn constrains the power of the event from above. This time is derived in section \\ref{sec:time} and drawn by the thick blue line on Figure \\ref{fig:totEvst}."
    },
    "1107/1107.1498_arXiv.txt": {
        "abstract": "Observations of the clustering of galaxies can provide useful information about the distribution of dark matter in the Universe. In order to extract accurate cosmological parameters from galaxy surveys, it is important to understand how the distribution of galaxies is biased with respect to the matter distribution. The large-scale bias of galaxies can be quantified either by directly measuring the large-scale ($\\lambda\\gta60~\\mpch$) power spectrum of galaxies or by modeling the halo occupation distribution of galaxies using their clustering on small scales ($\\lambda\\lta30~\\mpch$). We compare the luminosity dependence of the galaxy bias (both the shape and the normalization) obtained by these methods and check for consistency. Our comparison reveals that the bias of galaxies obtained by the small scale clustering measurements is systematically larger than that obtained from the large scale power spectrum methods. We also find systematic discrepancies in the shape of the galaxy bias-luminosity relation. We comment on the origin and possible consequences of these discrepancies which had remained unnoticed thus far. ",
        "introduction": "\\label{sec:intro}  Large scale galaxy redshift surveys provide a unique window to probe the distribution of dark matter. The abundance and clustering of galaxies contain enormous information about the dark matter distribution in the Universe and hence various cosmological parameters \\citep[e.g., see][]{vdb2003,vdb2007,Tinker2005,Cacciato2009}. However, in order to use galaxies as cosmological probes, an accurate knowledge of how the galaxy distribution is biased with respect to the dark matter distribution is essential.  The Sloan Digital Sky Survey \\citep[][SDSS hereafter]{York2000} has imaged nearly a quarter of the sky ($\\sim$ 8000 sq. deg.) and obtained accurate spectroscopic redshifts for more than 0.9 million galaxies over a span of 8 years of its operation \\citep{Abazajian2009}. The availability of the three dimensional positional information for such a large sample of galaxies has enabled an accurate measurement of the galaxy power spectrum over more than an order of magnitude in wavelength range \\citep[][T04 hereafter]{Tegmark2004}. The shape of the power spectrum of galaxies has been used in combination with the results from the cosmic microwave background to put stringent constraints on cosmological parameters such as the matter density parameter, $\\Omega_m$, the vacuum energy density parameter, $\\Omega_\\Lambda$, its equation of state, $w_0$, and constraints on quantities relevant to particle physics such as the sum of masses of neutrinos \\citep[e.g., see][]{Tegmark2004b,Seljak2005,Reid2010}.  The cosmological parameter $\\sigma_8$ describes the variance of the matter density distribution smoothed on a scale of $8~\\mpch$ (comoving) at an early time in the Universe extrapolated to today using linear perturbation theory. The quantity $\\sigma_8^2$ is proportional to the amplitude of the matter power spectrum. As galaxies are biased tracers of the matter distribution, observations of galaxy power spectra on large scales ($\\lambda\\gta60~\\mpch$) can only constrain the product of galaxy bias and the parameter $\\sigma_8$.  This degeneracy between $\\sigma_8$ and galaxy bias can be broken by information from the clustering of galaxies (or by the clustering of the dark matter around galaxies probed by galaxy-galaxy lensing) on small scales ($\\lambda\\lta30~\\mpch$). However the clustering on these scales is inherently difficult to model. The halo model framework has been routinely used to carry out such modeling \\citep[e.g., see][]{Cooray2002}. The halo occupation distribution (HOD) of galaxies constrained by these models using the small-scale clustering predicts the large-scale bias of galaxies whose product with $\\sigma_8$ should agree with the power spectrum measurements modulo the systematic uncertainties in either of the methods.  In this paper, we compile and critically compare the measurements of the luminosity dependence of galaxy bias on large scales obtained using the measurements of the galaxy power spectrum and those obtained by modeling the clustering of galaxies on small scales. This allows us to gauge the systematic differences in the galaxy bias obtained from these different methods. In Section \\ref{sec:diff}, we briefly describe the different methods used to obtain galaxy bias and compare the results after accounting for the differences in the cosmological parameters used by each of the measurements. In Section~\\ref{sec:discuss}, we discuss the various sources of systematics in these methods and summarize our findings. ",
        "conclusions": "\\label{sec:discuss}  We have compiled measurements of the luminosity dependence of the galaxy bias obtained using various methods and compared the shape and the normalization of the galaxy bias. We find that the bias of $L_*$ galaxies with magnitude $^{0.1}M_r=-20.44$ is at best constrained to an accuracy of order 10-15\\% (considering both statistical and systematic errors). We find that the galaxy bias obtained from large-scale power spectrum methods is systematically lower than that obtained from the small-scale clustering measurements. The discrepancy increases as a function of galaxy luminosity. We have also shown that this systematic discrepancy is not a result of assuming cosmological parameters which are different from the concordance cosmological parameters. We briefly comment on the plausible origins of the discrepancy and its consequences.  \\subsection{Typographical errors} We first examine whether typographical errors in any of the manuscripts could have caused this discrepancy. The number which has the potential to raise a few eyebrows is the value of $M_*$ used by T04. The luminosity function has been measured with a great accuracy using SDSS data by \\citet{Blanton2003} who obtain $M_*=-20.44\\pm0.01$, instead of $-20.83$ used by T04.  We carried out the following exercise to verify that the $M_*$ to be used in Eq.~\\ref{eq:teg04} from T04 is indeed $-20.83$. Using the values of $b/b_*$ obtained by T04, our own fitting routines and assuming $M_*=-20.83$, we obtain the values of parameters $(A,B,C)$ in Eq.\\ref{eq:teg04} to be $(0.856,0.144,-0.038)$ consistent with T04. Instead, if we use $M_*=-20.44$, we obtain $(A,B,C)=(0.928,0.072,-0.068)$ and a fit which is significantly worse than the former. This rules out the possibility of typographical error by T04 while quoting the value $M_*$.  The value of $M_*=-20.83$ used by T04 \\citep[and also later by][]{Swanson2008} appears to be taken from \\citet{Blanton2001} where the luminosity function of SDSS galaxies was measured from early commissioning data (M. Blanton 2011, private communication). This value of $M_*$ was later rectified in \\citet{Blanton2003} by including a model for the passive evolution of galaxies and using a fitting function more flexible than the Schechter function. It appears that T04 failed to incorporate this change in their analysis.  However, this does not invalidate their power spectrum measurements. The power spectrum presented in T04 should be interpreted as that of galaxies with $^{0.1}M_r=-20.83$ as we have done in this paper, instead of that of $L_*$ galaxies.  \\subsection{Systematic issues: Power spectrum method} The large scale power spectrum measurements from T04 may also have some systematic uncertainties. To calculate the three dimensional power spectrum of galaxies, T04 had to remove the small scale redshift space distortions (finger-of-god effect, FOG hereafter) induced by the peculiar velocities of the galaxies. Their method involved using the friends-of-friends algorithm \\citep{Huchra1982} in redshift space to identify overdense regions (the overdensity $\\delta_c$ is a free parameter of the algorithm, T04 used $\\delta_c=200$ as the fiducial value) in the galaxy distribution. The FOG removal algorithm then measures the dispersion of galaxy positions about the center of the identified overdense region in both the radial and transverse directions and then compresses the positions of galaxies radially until the dispersions are equal in both the transverse and radial directions. T04 used an anisotropic metric to measure distances between galaxies to account for the anisotropic stretching of galaxies in the line-of-sight direction. Two galaxies were marked as friends of each other if their separations satisfy \\begin{eqnarray} \\left[ \\left(\\frac{r_{||}}{10}\\right)^2 + r_{\\perp}^2 \\right]^{1/2} &\\leq& \\left[ \\frac{4}{3}\\pi(1+\\delta_c)\\bar{n} + r_{\\rm \\perp max}^{-3} \\right]^{-1/3} \\\\ &\\leq& \\left[ \\frac{4}{3}\\pi(1+\\delta_c) +\\left(\\frac{r_{\\rm \\perp max}}{\\bar{l}}\\right)^{-3}\\right]^{-1/3} \\bar{l} \\end{eqnarray} where, $\\bar{n}$ is the mean number density of galaxies, $\\bar{l}=\\bar{n}^{-1/3}$ is the mean galaxy separation, $r_\\perp$ is the projected distance, $r_{||}$ is the distance along the line-of-sight and $r_{\\rm \\perp max}$ was set to be equal to $5~\\mpch$ by T04. This implies that the value of the linking length in units of the mean galaxy separation when $\\delta_c=200$ is $\\approx 0.106$ (ignoring the $r_{\\rm \\perp max}$ term which is negligible compared to the $\\delta_c$ term).  For such a small value of linking length the overdensity of structures identified by FOF should be $\\sim 2000$, i.e. much larger than $200$ \\citep{More2011} which T04 aim to find and this can be a possible cause of systematics. \\footnote{A C++ code to calculate the overdensity of FOF haloes for any given cosmology and linking length parameter is available at \\href{http://kicp.uchicago.edu/~surhud/research/odcode.tgz}{http://kicp.uchicago.edu/$\\sim$surhud/research/odcode.tgz}.}  The measurements of the galaxy bias by \\citet{Swanson2008} based on a counts-in-cells analysis, which agree with T04, also applied the same algorithm for FOG removal.  It must be noted however that T04 included a check for systematics by changing the value of $\\delta_c$ to be lower than their fiducial value which causes the linking length to be larger than $0.106$ and did not find large systematic effects on the measured galaxy power spectrum compared to the statistical errorbars (see Figures $20$ and $21$ in T04).  \\subsection{Systematic issues: Small scale clustering} Modeling the small scale clustering of galaxies is inherently difficult as it requires knowledge of the highly uncertain physics of galaxy formation. The halo model conveniently by-passes this issue by using a statistical description of how galaxies populate halos as opposed to a full fledged model for galaxy formation physics. However, such a modeling has its own issues. Its unclear how differences in the assumed parametric forms for the HOD affect the modeling results. On small scales, pairs of galaxies in the same halo dominate the clustering signal. An accurate knowledge of the number density distribution of galaxies in a single halo and the satellite fraction (fraction of galaxies of a given luminosity which are satellites) are required to model this signal \\citep[e.g., see][]{Mandelbaum2006, Yang2008, More2009}. The occupation number of satellites in halos of a given mass is often assumed to be distributed in a Poisson manner, which is perhaps a questionable assumption \\citep[e.g., see][]{Wetzel2010}. The transition region where the clustering of galaxies is equally composed of pairs from the same halo and those from separate halos has also been notoriously difficult to model. An accurate knowledge of the scale dependence of the bias and an accurate treatment of halo exclusion is central to model the clustering in this region \\citep{Tinker2005, Smith2011}. Given all these complications, it is important to test for consistency of the small scale clustering results with the independent measurement from the large-scale power spectrum of galaxies, especially as precision constraints on cosmological parameters are being obtained from the modeling of small-scale clustering \\citep[see e.g.,][]{Tinker2011}.  Another potential systematic effect concerns the range of integration used along the line-of-sight, $\\pi_{\\rm max}$, to project the 3-dimensional correlation function to obtain $w_p(r_p)$. In Figure $1$ of \\citet{Norberg2009}, the authors show that the projection of the correlation function by integrating in redshift space upto $\\pi_{\\rm max}=64~\\mpch$ gives results that are systematically larger compared to results when the integration is carried out in real space. The difference between these results is as large as $10\\%$ on scales of $r_p=10~\\mpch$ with differences which increase with scale and grow as large as $30\\%$ at scales of $30~\\mpch$. The projection is supposed to get rid of the redshift distortions but only when the integration is done along the entire line-of-sight. However, when the integration is done only out to a certain maximum value ($\\pi_{\\rm max}=60~\\mpch$ is a commonly used value), the large scale redshift distortions (known as the Kaiser effect) may still persist. This can cause an overestimate of the correlation function corresponding to an overestimate of the galaxy bias. The exact overestimate will depend upon the quantity \\begin{equation} \\beta = \\frac{1}{b}\\frac{\\drm \\ln D(z)}{ \\drm \\ln (1+z)^{-1} } \\,. \\end{equation} Our preliminary estimates of the effect indicate that it can reduce the discrepancy between the two measures of the galaxy bias but not entirely eliminate it. The galaxy bias measurements from Z10 when the Kaiser effect is included should be smaller than those shown in Figure~\\ref{fig1} by roughly 1 $\\sigma$. We intend to pursue the issue of estimating the magnitude of the Kaiser effect on the projected clustering measurements, its dependence on $\\pi_{\\rm max}$ and comparisons with N-body simulations in future work.  \\subsection{Possible implications} The luminosity dependence of the galaxy bias obtained by T04 was used in conjunction with HOD modeling of the galaxy-galaxy lensing results by \\citet{Seljak2005} to obtain constraints in the $b_*-\\sigma_8$ plane. It is unclear how the likelihood contours obtained by these authors will shift in the $b_*-\\sigma_8$ plane if they were to use the small scale clustering results instead. Future analyses which combine galaxy-galaxy lensing and clustering data should be able to shed more light on this issue. The luminosity dependence of the galaxy bias has also been used to obtain an effective halo mass for galaxies of a given luminosity \\citep[e.g., see][]{Hand2011}. One should be aware that the systematics highlighted in this paper will lead to large uncertainties in the halo masses especially at the bright end.  The LRG power spectrum measurements \\citep{Reid2010} use the $b(L)$ obtained by T04 to correct for the red-tilt of the galaxy power spectrum.  Systematic uncertainties in the shape of $b(L)$ shown in this paper can cause a residual scale-dependence in the galaxy power spectrum modifying its shape and systematically bias the estimates of the cosmological parameters derived from it. A consistent understanding of galaxy bias and the systematics associated with each of the methods is therefore crucial to obtain precision constraints on cosmological parameters from the galaxy distribution."
    },
    "1107/1107.3986_arXiv.txt": {
        "abstract": "We report on two fossil groups of galaxies at redshifts $z = 0.425$ and $0.372$ discovered in the Cosmic Evolution Survey ({\\it COSMOS\\/}) area. Selected as X-ray extended sources, they have total masses ($M_{200}$) equal to $1.9~(\\pm 0.41) \\times 10^{13}$ and $9.5~(\\pm 0.42) \\times 10^{13}~\\mathrm{M}_{\\sun}$, respectively, as obtained from a recent X-ray luminosity--mass scaling relation. The lower mass system appears isolated, whereas the other sits in a well-known large-scale structure (LSS) populated by 27 other X-ray emitting groups. The identification as fossil is based on the $i$-band photometry of all the galaxies with a photometric redshift consistent with that of the group at the $2 \\sigma$ confidence level and within a projected group-centric distance equal to $0.5 R_{200}$, and $i_{\\mathrm AB} \\le 22.5~\\mathrm{mag}$ limited spectroscopy. Both fossil groups exhibit high stellar-to-total mass ratios compared to all the X-ray selected groups of similar mass at $0.3 \\le z \\le 0.5$ in the {\\it COSMOS\\/}. At variance with the composite galaxy stellar mass functions (GSMFs) of similarly massive systems, both fossil group GSMFs are dominated by passively evolving galaxies down to $M^{stars} \\sim 10^{10}~\\mathrm{M}_{\\sun}$ (according to the galaxy broad-band spectral energy distributions). The relative lack of star-forming galaxies with $10^{10} \\le M^{stars} \\le 10^{11}~\\mathrm{M}_{\\sun}$ is confirmed by the galaxy distribution in the $b - r$ vs $i$ colour--magnitude diagram. Hence, the two fossil groups appear as more mature than the coeval, similarly massive groups. Their overall star formation activity ended rapidly after an accelerated build up of the total stellar mass; no significant infall of galaxies with $M^{stars} \\ge 10^{10}~\\mathrm{M}_{\\sun}$ took place in the last 3 to 6 Gyr. This similarity holds although the two fossil groups are embedded in two very different density environments of the LSS, which suggests that their galaxy populations were shaped by processes that do not depend on the LSS. However, their progenitors may do so. We discuss why the late merging of a compact group is favoured over the early assembly as a formation scenario for the isolated, low-mass fossil group. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.1384_arXiv.txt": {
        "abstract": "I investigate the suggestion that the Hyades moving group in the Solar neighbourhood is the result of a recent inner Lindblad resonance. I use dynamical ``torus'' models of the Galaxy to understand the expected distribution of solar neighbourhood stars in angle coordinates for phase-mixed models and models which include a resonant component. I show that attempts to find the signatures of resonances in angle coordinates are strongly influenced by selection effects, including rather subtle effects associated with the relationship between action and angle for stars at a given point. These effects mean that one can not use simple tests to determine whether substructures seen in the Solar neighbourhood are associated with any given resonance. ",
        "introduction": "\\label{sec:intro} Since \\cite{WD98} used proper motions and parallaxes obtained by the Hipparcos satellite \\citep{Hipparcos} to investigate the kinematics of the Solar neighbourhood, it has been clear that the local distribution function (\\df) is far from smooth. In particular, the distribution of stars in the $U,V$ plane\\footnote{Throughout this paper, in the Solar neighbourhood, velocities with respect to the Sun are described in terms of a component towards the Galactic Centre ($U$), a component in the direction of Galactic rotation ($V$), and a component perpendicular to the Galactic plane towards the north Galactic pole ($W$). Velocities with respect to the local standard of rest are described in terms of components in the same directions as $U$, $V$ and $W$ which are given the symbols $v_x$, $v_y$ and $v_z$ respectively} is dominated by a number of streams, all of which are thought to be dynamical in origin \\citep[e.g.][]{Faea05}.  Recently \\citet[henceforth S10]{Se10} argued that part of this substructure could be explained by a recent inner Lindblad resonance (ILR), a conclusion he supported with reference to the distribution in angle coordinates of stars observed by the Geneva-Copenhagen survey \\citep*[GCS:][]{GCS09}. This conclusion was supported by \\cite*{HaSePr11} who looked at stars in the Solar neighbourhood observed by the Radial Velocity Experiment \\citep[RAVE:][]{RAVE1_short} and the Sloan Digital Sky Survey \\citep[SDSS:][]{SDSS7_short}.  In this paper I reexamine S10's conclusion that the distribution of stars in the Solar neighbourhood show signs of an inner Lindblad resonance (ILR). I compare the GCS sample of stars in the Solar neighbourhood to a phase-mixed dynamical model. This allows me to separate selection effects from genuine substructure in the local \\df , and to develop some understanding of the impact of selection effects on the observed properties of any substructure that is found in the local \\df . I then explore the impact of selection effects on simple models of an ILR or an outer Lindblad resonance (OLR).  In Section~\\ref{sec:AAcoords} I discuss angle-action coordinates, how they might be used to determine the dynamical origin of observed kinematic substructure, and their relationship to kinematics in the Solar neighbourhood. In Section~\\ref{sec:num} I give numerical details of the assumptions made and the phase-mixed dynamical model considered. In Section~\\ref{sec:smooth} I explore the appearance in angle coordinates of the solar neighbourhood in a phase-mixed model, which I then use in Section~\\ref{sec:GCS} to interpret the distribution in angle coordinates of stars observed by the GCS. Section~\\ref{sec:resmod} discusses simple models which include a resonant component, and uses them to better understand the GCS data. ",
        "conclusions": "\\label{sec:conc} In this paper I have re-examined the distribution of stars in the Solar neighbourhood in angle coordinates following the claim by S10 that the Hyades moving group is related to an ILR. Using a dynamical ``torus'' model I showed the significant impact of selection effects associated with surveying a finite (small) volume upon the distribution of stars in angle coordinates taken from a phase-mixed model. Using models which contain resonant components in addition to a phase-mixed background I have demonstrated the important effects that the distribution of resonant stars in action have on the observed distribution in angle (again because of selection effects).  The distribution of the stars associated with the Hyades moving group in action (Figure~\\ref{fig:actions}) indicates that it is associated with some resonance between the radial oscillations and the azimuthal motion (i.e. follows a relation of the form of $l\\Omega_r+m\\Omega_\\phi$). It is also clear that the stars which make up the resonant component are constrained in $\\bolth$ in some way, otherwise there would be an overdensity at $\\theta_r\\sim1.5$ as well as $\\theta_r\\sim-1.5$. However, it is very difficult to determine what the values of $l$ and $m$ are, and thus which resonance is responsible. In action space the lines $l\\Omega_r(\\bolJ)+m\\Omega_\\phi(\\bolJ)=const$ are very similar, in the relevant range of $\\bolJ$, for different values of $l$ and $m$, and are sensitive to choice of the Galactic potential. The approach taken by S10, of looking for overdensities in the statistic $l\\theta_r+m\\theta_\\phi$, is also flawed because of the selection effects that have important influences on the distribution in $\\bolth$. I find that, contrary to the conclusions of S10, it not clear that the Hyades moving group is the result of an inner Lindblad resonance.  It may be possible to use careful modeling and analysis, which takes into account selection effects and the distributions in $\\bolth$ and $\\bolJ$ simultaneously, to determine what type of resonance is responsible for the Hyades moving group from these data. However this task is made even more difficult by uncertainty about the Galactic potential which significantly affects the gradient of resonance lines in action space.  It seems likely the problems of selection effects in studies of this kind will prove extremely challenging, if not intractable, while we are dealing with survey volumes corresponding to the relatively small radii ($\\sim200\\pc$) associated with the GCS. To go beyond this with similar accuracy (i.e.~$\\sim1\\kms$ velocity uncertainty) requires distance and proper measurements more accurate than those currently available -- the RAVE and SDSS data used by \\cite{HaSePr11} has significantly larger uncertainties even though they select stars which lie in the same volume as the GCS stars. Gaia \\citep{GAIA01} is expected to produce 3-dimensional velocity measurements accurate to $\\sim1\\kms$ for some stars up to $\\sim3\\kpc$ from the Sun, corresponding to a range in Galactocentric $\\phi$ of order $0.7$ radians. This will dramatically reduce the impact of selection effects on the observed distributions in angle-action coordinates, and should make it relatively straightforward to pick out structures in angle space that are unambiguously of the kind predicted by S10. This process is likely to both be guided by and affect the determination of the gravitational potential of the Galaxy (an essential product of the Gaia mission).  Throughout this study I have assumed that the uncertainties on the positions and velocities of the stars observed by the GCS can be ignored, because they are too small to have any significant effect. It is worth noting that the relationship between the uncertainty in $\\bolx,\\bolv$ and that in $\\bolth,\\bolJ$ is not entirely straightforward, and that a small uncertainty in velocity does not necessarily correspond to a small uncertainty in $\\bolth$. As inspection of Figure~\\ref{fig:UVaxes} suggests, the relationship between an uncertainty in velocity and that in $\\theta_r$ is heavily dependent on $J_r$, with smaller $J_r$ corresponding to a greater uncertainty in $\\theta_r$ for given uncertainty in $\\bolv$. Indeed as $J_r\\rightarrow0$ a negligible uncertainty in $\\bolv$ corresponds to an uncertainty of $2\\pi$ in $\\theta_r$. The Hyades overdensity does not extend to $J_r=0$ (as can be seen in Figure~\\ref{fig:UV_GCS}), so the assumption that the uncertainty in $\\theta_r$ can be ignored is reasonable for this study. However, future efforts to understand resonances of this type is likely to require appropriate treatment of these uncertainties."
    },
    "1107/1107.1517_arXiv.txt": {
        "abstract": "We present results from daily radio continuum observations of the \\boo field as part of the Pi GHz Sky Survey (PiGSS).  These results are part of a systematic and unbiased campaign to characterize variable and transient sources in the radio sky. The observations include 78 individual epochs distributed over 5 months at a radio frequency of 3.1 GHz with a median RMS image noise in each epoch of 2.8 mJy. We produce 5 monthly images with a median RMS of 0.6 mJy. No transient radio sources are detected in the daily or monthly images. At 15 mJy, we set an upper limit ($2\\sigma$) to the surface density of 1-day radio transients at 0.025 deg$^{-2}$.  At 5 mJy, we set an upper limit ($2\\sigma$) to the surface density of 1-month radio transients at 0.18 deg$^{-2}$.  We also produce light curves for 425 sources and explore the variability properties of these sources.  Approximately 20\\% of the sources exhibit some variability on daily and monthly time scales.  The maximum RMS fractional modulations on the 1 day and 1 month time scales for sources brighter than 10 mJy are 2 and 0.5, respectively.  The probability of a daily fluctuation for all sources and all epochs by a factor of 10 is less than $10^{-4}$. We compare the radio to mid-infrared variability for sources in the field and find no correlation.  Finally, we apply the statistics of transient and variable populations to constrain models for a variety of source classes. ",
        "introduction": "The temporal properties of the radio wavelength sky are poorly understood. Much of what is known results from follow-up of events discovered at optical, X-ray, or $\\gamma$-ray wavelengths such as supernovae, gamma-ray burst, and X-ray binaries.  An important exception is the sub-second domain, which has been systematically, if incompletely, explored, leading to the discovery of pulsars \\citep{1969Natur.224..472H}, rotating radio transients \\citep{2006Nature...mcl}, and other unexplained phenomena \\citep{2007Sci...318..777L,2009GeoRL..3613202R,2010MNRAS.402..855B,2011arXiv1104.2727K}.  Longer-time scale searches, which tend to be sensitive to synchrotron sources, are rarer and the sky coverage is less complete.  Systematic searches on long timescales for radio transients (RTs) are undergoing a renaissance that is defining the available parameter space.  These searches are motivated by the specific promise of orphan gamma-ray burst afterglows \\citep{2008MNRAS.390..675R}, radio supernovae \\citep{2009A&A...499L..17B}, tidally-disrupted stars around massive black holes \\citep{2011arXiv1102.1429G,2011arXiv1103.4328B}, intrinsic AGN activity such as found in III Zw 2 \\citep{1999ApJ...514L..17F}, extreme scattering events \\citep{1987Natur.326..675F}, stellar events \\citep{2007ApJ...663L..25H}, and counterparts to gravity-wave events \\citep{2011arXiv1102.1020N}.  Additionally, these searches are motivated by the desire to fully characterize the radio sky and discover the unexpected and unpredicted.  Archival searches have proven fruitful in this search, uncovering a number of transients that are not easily explained with known source classes \\citep{2002ApJ...576..923L,2006ApJ...639..331G,2007ApJ...666..346B,2010AJ....140..157B,2011ApJ...728L..14B,2011MNRAS.412..634B,2011arXiv1103.0511B}. These archival surveys have been powerful for the volume probed but the absence of the opportunity for real-time follow-up has significantly limited their value in characterizing the source population. Synoptic transit surveys have also been carried out \\citep{2008NewA...13..519K,2009AJ....138..787M}. In particular, a number of new variable and transient sources without deep radio or optical counterparts have been found.  \\citet{2010ApJ...711..517O} argue that these may be bursts from galactic neutron stars; other evidence suggests that they could be strong bursts from low-mass flare stars.  The success of archival surveys has lead to a growing number of unbiased surveys including those carried out by the Allen Telescope Array \\citep[ATA;][]{2009IEEEP..97.1438W}.  The ATATS survey places limits on rare, high flux density transients and variables \\citep{2010ApJ...719...45C,2011ApJ...731...34C}. A recent VLA/EVLA survey \\citep{2011arXiv1103.3010O} has now provided strong limits on the radio transient and variable population at 5 GHz.  These authors also summarize existing RT surveys at frequencies $>0.8$ GHz.  Long wavelength transients were searched for with the Long Wavelength Array \\citep{2010AJ....140.1995L}. These archival and unbiased surveys join targeted surveys of specific source classes in characterizing the radio transient sky. For example, the MASIV survey \\citep{2003AJ....126.1699L,2008ApJ...689..108L} and the University of Michigan monitoring program \\citep{1992ApJ...396..469H} provide characterization of bright, flat-spectrum radio sources, which are known to have the strongest variability amplitudes.  The Pi GHz Sky Survey (PiGSS) is a key project of the ATA. In \\citet{2010ApJ...725.1792B}, hereafter Paper I, we presented the survey description, methods in observation, data reduction, and source identification, and initial results on transient statistics. Briefly, PiGSS is a 3.1 GHz survey of a large area of high galactic latitude sky conducted synoptically with 100\\arcsec\\ resolution and milli-Jansky sensitivity.  The PiGSS emphasis is on two surveys jointly conducted: 1) $10^4$ deg$^2$ of the sky with $b > 30^\\circ$; and 2) individual $\\sim 11$ deg$^2$ fields that are observed repeatedly. Paper I presented results from integrated results for a 11-deg$^2$ field in the \\boo constellation associated with the NOAO Deep Wide Field Survey \\citep[NDWFS;][]{2004AAS...205.8106J,2009ApJ...701..428A}. Numerous sensitive, large area surveys have been carried out for the \\boo field, providing an important baseline for comparison with PiGSS results.  We found no new radio sources that we could convincingly identify as RTs.  This sets an upper limit on the RT surface density at the milliJansky flux limit.  In this paper, we present daily results from the \\boo field.  These consist of 78 individual epochs over five months. These observations provide a very sensitive probe of transient statistics.  They also provide one of the first unbiased searches for variability in a sample of mJy brightness sources.  In \\S \\ref{sec:obs}, we briefly review the observations and data analysis and present the full catalog of light curves.  In \\S \\ref{sec:transients}, we present results on transient sources.  In \\S \\ref{sec:var}, we characterize the variability of individual sources.  In \\S \\ref{sec:conc}, we summarize our results. ",
        "conclusions": "}  We report on daily observations of an 11-deg$^2$ region in the \\boo constellation with the ATA at 3.1 GHz as part of the PiGSS project.  The 78 epochs of observation provide a systematic look at daily and monthly variations in the radio sky.  We find no radio transients.  There is a single event that was detected in the deep image without a counterpart in pre-existing surveys such as NVSS that, if real, has a duration $\\gsim 4$ months.  We place a new constraint on the surface density of 1-day duration transients with $S > 15$ mJy of 0.025 deg$^{-2}$ and monthly duration transients with $S > 5$ mJy of 0.18 deg$^{-2}$.  These limits are consistent with previous efforts with the ATA and with other instruments.  We also explored variability of the 425 radio sources in the deep field through a variety of statistics. Largely, these sources show weak or no variability with RMS fractional modulations less than 1 for sources brighter than 10 mJy.  20\\% of the sources have statistically significant variability. The combination of 425 sources and 78 epochs provides over $3 \\times 10^4$ flux density samples from which we can place limits on rare, large amplitude fluctuations.  The largest amplitude fluctuations ($>10 \\times$) occur fewer than 1 in $10^4$ source-days.  EVLA characterization of a subset of these sources is important for determining the effect of systematic errors in some of the larger variations reported here.  While these statistics are of the most use for broadly characterizing the population, certain classes of objects are better identified with matched filtering techniques.  For instance, extreme scattering events and radio supernovae have reasonably well-defined and parametrizable light curves which could be used to explore these data sets in greater depth.  Future explorations of data of this kind should include matched filter analysis.  The PiGSS results in the \\boo field represent approximately 20\\% of the data from the PiGSS project and were obtained in the early phases of routine telescope operations. Subsequent analysis of higher quality daily observations in the Lockman Hole, ELAIS, and the Coma cluster will probe significantly deeper into the radio transient and variable population.  PiGSS observations will also characterize variability on time scales of 1 year for $\\sim 10^5$ radio sources in the North Galactic Cap.  The limits we have placed are useful constraints on the population of active high energy objects in the Universe.  These limits explore the populations of tidal disruption events, radio supernovae, orphan gamma-ray burst afterglows, radio counterparts to GW sources, extreme scattering events, intra-day variable sources, and flare stars."
    },
    "1107/1107.0277_arXiv.txt": {
        "abstract": "In this study, we present photometric and spectroscopic variations of the extremely small mass ratio ($q\\simeq 0.1$) late-type contact binary system \\astrobj{V1191 Cyg}. The parameters for the hot and cooler companions have been determined as $M_\\textrm{h}$ = 0.13 (1) $M_{\\odot}$, $M_\\textrm{c}$ = 1.29 (8) $M_{\\odot}$, $R_\\textrm{h}$ = 0.52 (15) $R_{\\odot}$, $R_\\textrm{c}$ = 1.31 (18) $R_{\\odot}$, $L_\\textrm{h}$ = 0.46 (25) $L_{\\odot}$, $L_\\textrm{c}$ = 2.71 (80) $L_{\\odot}$, the separation of the components is $a$= 2.20(8) $R_{\\odot}$  and the distance of the system is estimated as 278(31) pc. Analyses of the times of minima indicates a period increase of  $\\frac{dP}{dt}=1.3(1)\\times 10^{-6}$ days/yr that reveals a very high mass transfer rate of $\\frac{dM}{dt}=2.0(4)\\times 10^{-7}$$M_{\\odot}$/yr from the less massive component to the more massive one. New observations show that the depths of the minima of the light curve have been interchanged. ",
        "introduction": "\\label{intro} Modelling the evolution of late-type contact binary systems (LTCBs) is highly complicated. Yakut \\& Eggleton (2005) (hereafter YE05) modelled late-type close binary systems by assuming conservative and non-conservative cases. The authors proposed that a detached configuration can evolve into a semi-detached and then to a contact configuration. The mass ratio of the components reverses through evolution. Therefore, currently observed low mass star could be initially the more massive one in a binary system. The mass ratio of the components, which is related to the angular momentum loss and mass transfer, is one of the crucial parameter in  the evolution of close binary systems. Apart from \\astrobj{V1191 Cyg}, there exist rare well-known LTCBs with a small mass ratio (e.g. \\astrobj{SX Crv}, \\astrobj{V870 Ara}, \\astrobj{FG Hya}, \\astrobj{$\\epsilon$ CrA}, \\astrobj{CK Boo}, \\astrobj{FP Boo}, \\astrobj{DN Boo}, \\astrobj{AH Cnc}; see YE05, {\\c S}enavc{\\i} et al. 2008, Yakut et al. 2009).  \\astrobj{\\astrobj{V1191 Cyg}} was classified as a W UMa type system by Mayer in 1965. Many years later, the first photometric study of the system was made by Pribulla et al. (2005a) and they derived the geometric elements of the binary system without a spectroscopic study. The authors gave the mass ratio of the components as 0.09 and reported orbital period increase of $4.2\\times10^{-6}$ yr$^{-1}$. Rucinski et al. (2008) studied the radial velocity curve of \\astrobj{V1191 Cyg} and determined the spectroscopic mass ratio, orbital period, radial velocity amplitudes, and the velocities of the center of mass of the system. However, the authors also mentioned that the system is indeed a difficult spectroscopic target because of its faintness. ",
        "conclusions": "B, V, and R light curves of the system have been solved simultaneously with the spectroscopic study of Rucinski et al. (2008). We derived the orbital and the physical parameters of the components. Using obtained parameters we could estimate the filling factor ($f=\\frac{\\Omega_{{in}}-\\Omega}{\\Omega_{{in}}-\\Omega_{{out}}}$) from the inner (${\\Omega_{{in}}}$) to the outer critical surface (${\\Omega_{{out}}}$) as 0.74. The distance of the system to the Sun estimated as 278 pc by using observed parameters (Table~\\ref{tab:V1191Cyg:par}). This is the first distance value given for \\astrobj{V1191 Cyg}. During the calculations, effective temperature and absolute magnitude of the Sun are taken as 5780~K and 4.75 mag, respectively. We compared our result with the well-known contact binaries using YE05 tables and figures. The physical parameters obtained from our analysis of \\astrobj{V1191 Cyg}  seem to be in a good agreement with the well-known LTCBs. Like the other W-subtype secondaries, the less massive component of \\astrobj{V1191 Cyg} is overluminous and oversized in YE05 diagrams. While the secondary component of the system appears to be below the ZAMS, the massive and the cooler component is situated close to the TAMS. This phenomenon was discussed in detail by YE05, Kalomeni et al. (2007),  Stepien (2006), and Gazeas \\& Niarchos (2006).   Light variations at maxima and minima depths have been detected for the first time. The light curves of Pribula et al. (2005a) indicate that the occultation minimum is deeper than the transit by $\\sim 0.01$ mag in R band. However, light curves presented in this study (Fig.~\\ref{Fig:V1191Cyg:LC}a) show that the deeper minimum is the transit and the shallower one is the occultation. The occultation minimum is shallower by $\\sim 0.04$ mag in R band.   There are two subgroups of A- and W- types for LTCBs according to Binnendijk (1970) classification. If the more massive component (the hotter one) is eclipsed during the deeper minimum, then we are dealing with an A-type system. If the less massive star is the hotter one, the system is W-type. \\astrobj{FG Hya} (Qian \\& Yang, 2005) and \\astrobj{EM Psc (Qian} et al., 2008) show an interchange between this two subgroups which is consistent with the model of YE05. The light curves of Pribula et al. (2005a) show that \\astrobj{V1191 Cyg} is a W-type system whereas our observations show that the minima of the light curve have been interchanged. Apart from the long period evolution (e.g. heat transfer between components) of LTCB systems, short period variation (e.g. efficiently stellar activity) can affect this phenomena. \\astrobj{V1191 Cyg} shows similar characteristic properties of \\astrobj{FG Hya} and \\astrobj{EM Psc}.   \\astrobj{V1191 Cyg} with binaries having small mass ratio and short orbital period form a small but important sub-group of binary zoo. Short period and small mass ratio in contact binaries play a crucial role to explain astrophysical phenomena that are not well understood yet (i.e. merge of close/contact binaries, blue stragglers and the formation of \\astrobj{FK Com} stars) (YE05; Arbutina, 2009; Jiang et al., 2010). It is generally believed that these kinds of systems can be the progenitors of the \\astrobj{FK Com} systems and blue stragglers.   \\begin{table*} \\begin{center} \\caption{Absolute parameters of V1191 Cyg. The standard errors 1$\\sigma$ in the last digit are given in parentheses. H and C stand for the hot and the cooler components, respectively} \\label{tab:V1191Cyg:par} \\begin{tabular}{llll} \\hline Parameter                                        &Unit                      & H           & C   \\\\ \\hline Mass (M)                                         &$\\rm{M_{\\odot}}$      & $0.13(1)$            & $1.29(8)$      \\\\ Radius (R)                                       &$\\rm{R_{\\odot}}$      & $0.52(15)$           & $1.31(18)$      \\\\ Temperature (T$_{\\rm eff}$)                      &$\\rm{K}$              & $6610(200)$              & $6500$    \\\\ Luminosity (L)                                   &$\\rm{L_{\\odot}} $     & $0.46(8)$           & $2.71(44)$      \\\\ Surface gravity ($\\log g$)                       &$\\rm{cms^{-2}} $      & $4.12$             & $4.31$      \\\\ Absolute magnitude (M$_V$)                       &mag                   & 5.73              & 3.82           \\\\ Period change rate ($\\dot{P}$)                   &d/yr                 &~~~~~~~~~~~~~$1.3(1)\\times10^{-6}$ &      \\\\ Mass transfer ratio ($\\dot{M}$)                  &M$_\\odot$/yr         &~~~~~~~~~~~~~$2.0(4)\\times10^{-7}$ &      \\\\ Distance (d)                                     &pc                    &~~~~~~~~~~~~~~~~~278(31) &      \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}"
    },
    "1107/1107.4986_arXiv.txt": {
        "abstract": "We have identified a star in the WASP archive photometry with an unusual lightcurve due to the total eclipse of a small, hot star by an apparently normal A-type star and with an orbital period of only 0.668\\,d. From an analysis of the WASP lightcurve together with V-band and I$_{\\rm C}$-band photometry of the eclipse and a spectroscopic orbit for the A-type star we estimate that the companion star has a mass of $0.23\\pm0.03$\\Msolar\\ and a radius of $0.33\\pm0.01$\\Rsolar, assuming that the A-type star is a main-sequence star with the metalicity appropriate for a thick-disk star. The effective temperature of the companion is $13\\,400\\pm 1\\,200$K from which we infer a luminosity of $3\\pm1$\\Lsolar. From a comparison of these parameters to various models we conclude that the companion is most likely to be the remnant of a red giant star that has been very recently stripped of its outer layers by mass transfer onto the A-type star. In this scenario, the companion is currently in a shell hydrogen-burning phase of its evolution, evolving at nearly constant luminosity to hotter  effective temperatures prior to ceasing hydrogen burning and fading to become a low-mass white dwarf composed of helium (He-WD). The system will then resemble the pre-He-WD/He-WD companions to A-type and B-type stars recently identified from their Kepler satellite lightcurves (KOI-74, KOI-81 and KIC~10657664). This newly discovered binary offers the opportunity to study the evolution of a stripped red giant star through the pre-He-WD stage in great detail. ",
        "introduction": "Wide-area surveys for transiting extra-solar planets such as WASP (Wide Angle Search for Planets, \\citealt{2006PASP..118.1407P}),  HATnet \\citep{2004PASP..116..266B}, XO \\citep{2005PASP..117..783M} and TrES \\citep{2006AAS...20922602O} provide high cadence photometry for millions of bright stars across a large fraction of the sky. This provides the opportunity to find and study many new examples of known classes of variable star, e.g., eclipsing brown dwarf binary systems \\citep{ 2011ApJ...726L..19A}, double-mode RR~Lyr stars \\citep{2010IBVS.5955....1W}, W~UMa stars \\citep{2011A&A...528A..90N}, young solar-type stars \\citep{2011arXiv1104.2986M} and cataclysmic variable stars \\citep{2011JAVSO.tmp..140W}. New discoveries will certainly be made now that much of the data from these surveys is becoming widely available \\citep{2010A&A...520L..10B}. The photometric precision achieved by these surveys with modest equipment ($\\la 0.01$\\,magnitudes at V$\\approx$12) is impressive, but cannot compete with the micro-magnitude photometry acheived from space by surveys such as CoRoT \\citep{2006ESASP1306...33B} and Kepler \\citep{2009IAUS..253..289B}. Photometry with this precision has made  it possible to identify new types of variable star that are difficult or impossible to study from the ground, e.g., triply eclipsing binary stars \\citep{ 2011Sci...331..562C}, stars with tidally excited pulsations \\citep{ 2011arXiv1102.1730W}, subdwarfs with white dwarf companions \\citep{2011MNRAS.410.1787B},  and white dwarf companions to early-type main sequence stars \\citep{2010ApJ...713L.150R, 2010ApJ...715...51V}.  One advantage that ground-based surveys currently have over space-based surveys is that they cover a much larger fraction of the sky. This makes is possible in some cases to discover rare or extreme examples of these new classes of variable star. In the case of the white-dwarf companions to early-type stars KOI-74 and KOI-81, these were identified from the eclipses and transits in the Kepler lightcurves, even though these features are  much less than 1\\% deep \\citep{2010ApJ...713L.150R, 2010ApJ...715...51V}. These low-mass white dwarfs have an unsual evolutionary history, but it is difficult to study them in detail because they are very much fainter than their companions. Ground-based surveys offer the opportunity to discover similar eclipsing binary systems that are more favourable for detailed follow-up observations.  Low mass white dwarf stars ($M<0.4\\Msolar$) are the product of binary star evolution \\citep{1993PASP..105.1373I, 1995MNRAS.275..828M}. They are the result of mass transfer from a red giant onto a companion star when the giant has a small degenerate helium core. There are several possible outcomes  from this mass transfer depending on the mass ratio of the binary and the type of companion star. If the companion is a neutron star then the mass transfer is likely to stable so the binary can go on to become a low mass X-ray binary (LMXB) containing a millisecond pulsar. Several millisecond radio pulsars are observed to have low mass white dwarf (LMWD) companions \\citep{2008LRR....11....8L}. Many new LMWDs have recently been identified in the Sloan Digital Sky Survey \\citep{2007ApJ...660.1451K} and from proper motion surveys \\citep{2009A&A...506L..25K}.  Searches for radio pulsar companions to these LMWDs have so far found nothing, suggesting that the majority of these LMWDs have white dwarf companions \\citep{2009ApJ...697..283A}. LMWD can also be produced by mass transfer from a red giant onto a main sequence star, either rapidly through unstable common-envelope envolution or after a longer-lived ``Algol'' phase of stable mass transfer \\citep{1969A&A.....1..167R, 1970A&A.....6..309G, 2004A&A...419.1057W, 1993PASP..105.1373I, 2003MNRAS.341..662C, 2001ApJ...552..664N}.  The evolution of LMWDs is expected to be very different from more massive white dwarfs.  It is expected that once mass transfer stops the LMWD will have a thick layer of hydrogen surrounding the degenerate helium core which leads to steady hydrogen shell burning via the p-p chain and can lead to unstable phases of CNO burning for LMWD in the mass range 0.2\\,--\\,0.3\\Msolar\\ \\citep{1999A&A...350...89D}.  These hydrogen shell flashes lead to mixing between the inner and outer layers, producing a hydrogen deficient surface composition for the LMWD. Models that include hydrogen shell burning provide a much better match between the ``cooling age'' of the LMWD and the ``spin-down age'' of the millisecond pulsars in LMXBs. The number of hydrogen shell flashes depends strongly on the mass of the hydrogen layer that remains on the surface of the LMWD \\citep{2000MNRAS.316...84S}, which in-turn depends strongly on the details of the mass loss from the red giant \\citep{2002ApJ...565.1107P}. Details such as the treatment of diffusion can also lead to large differences in the predictions between different models \\citep{2001MNRAS.323..471A}.  In this paper we present the discovery from WASP photometry of an eclipsing binary star that is related to KOI-74, KOI-81 but that has much deeper eclipses, i.e. the companion to the early-type star is much brighter and larger than the white-dwarf companions to these Kepler discoveries. We present the data used to identify this new eclipsing binary star and the follow-up photometry and spectroscopy we have obtained; we analyse these data to determine the masses, radii and luminosities of the stars; and we outline how our discovery of this bright pre-white dwarf companion to an early-type star makes it possible to study the formation of an LMWD in detail.   \\begin{table} \\caption{Catalogue photometry and astrometry of \\J. The GALEX FUV and NUV fluxes are given as AB magnitudes f$_{\\rm AB}$ and n$_{\\rm AB}$, respectively. Sources are GALEX \\citep{2007ApJS..173..682M}, NOMAD \\citep{ 2004AAS...205.4815Z} PPMXL\\citep{2010AJ....139.2440R}, UCAC3 \\citep{ 2010AJ....139.2184Z},  2MASS \\citep{2006AJ....131.1163S}.} \\label{mags} \\begin{tabular}{@{}lrl} \\hline Parameter &\\multicolumn{1}{l}{Value} & Source \\\\ RA (J2000.0) & 02 47 43.38 & NOMAD \\\\ Dec (J2000.0) &  $-$25 15 49.3 & NOMAD \\\\ $\\mu_{\\alpha}$ & $23.2 \\pm 2.5$ mas/yr & PPMXL \\\\ & $25.9 \\pm 1.2$ mas/yr & UCAC3 \\\\ $\\mu_{\\delta}$ & $-5.4 \\pm 1.2$ mas/yr & PPMXL \\\\ & $-5.3 \\pm 1.0$ mas/yr & UCAC3 \\\\ f$_{\\rm AB}$  & $16.03 \\pm 0.02$ &  GALEX \\\\ n$_{\\rm AB}$  & $14.46 \\pm 0.01$ & GALEX \\\\ B & 12.13 & NOMAD \\\\ V & 12.44 & NOMAD \\\\ J & 11.85 & 2MASS \\\\ H & 11.81 & 2MASS \\\\ K$_s$ & 11.79 & 2MASS \\\\ \\hline \\end{tabular} \\end{table}  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{lcfit.eps} \\caption{Lightcurves of \\J\\ with lightcurve model fits (solid lines). From bottom-to-top: WASP (small points); SAAO 1.0-m I$_{\\rm C}$-band and V-band (filled circles); ASAS V-band (small crosses). \\label{lcfit}  } \\end{figure}  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{plotpar6.eps} \\caption{Distribution of selected lightcurve parameters from the bootstrap Monte Carlo simulation. \\label{plotpar6}  } \\end{figure}  \\begin{figure} \\includegraphics[angle=270,width=0.45\\textwidth]{spec.ps} \\caption{Flux calibrated and normalized GMOS-S spectrum of \\J\\ (lower spectrum) compared to the normalized, flux-calibrated and smoothed spectrum of the A6Vp star HD148898 offset by +0.5 units (upper spectrum). \\label{spec}  } \\end{figure}   \\begin{figure} \\includegraphics[width=0.45\\textwidth]{rvfit.eps} \\caption{Radial velocity measurements of J0247$-$25\\,B as a function of phase with a circular orbit fit (solid line). The spectrograph used is indicated as follows: filled circles -- EFOSC2; triangles -- ISIS; diamonds -- GMOS. \\label{rvfit}  } \\end{figure} ",
        "conclusions": "The star  1SWASP~J024743.37$-$251549.2  is an eclipsing binary star in which the precursor to a low mass white dwarf with a mass $\\approx 0.25$\\Msolar\\ is totally eclipsed by a larger, cooler star once every 0.6678\\,d. More detailed spectroscopy will be required to measure a precise masses for the stars. This will enable us to determine the nature of the larger star and to make detailed tests of models for the formation pre-He-WDs and low mass white dwarfs."
    },
    "1107/1107.5409_arXiv.txt": {
        "abstract": "{We present a multi-wavelength, high-resolution observational survey of the young binary system Haro 6-10 (GV~Tau, IRAS 04263+2426), which is harbouring one of the few known infrared companions. } {The primary goal of this project is to determine the physical and geometrical properties of the circumstellar and circumbinary material in the Haro~6-10 system. } {High-resolution optical (HST/WFPC2) and near-infrared (VLT/NACO) images in different bands were analysed to investigate the large-scale structures of the material around the binary. Mid-infrared interferometry (VLTI/MIDI) and spectroscopy (TIMMI2 at the 3.6m ESO telescope) were carried out to determine the structure and optical depth of the circumstellar material around the individual components. } {The multi-wavelength observations suggest that both components of the binary system Haro~6-10 are embedded in a common envelope. The measured extinction indicates a dust composition of the envelope similar to that of the interstellar medium. Each component of the system has a circumstellar disc-like structure typical of young stars. The discs are highly misaligned: the northern component is seen almost edge-on and the southern component is an almost face-on disc. % } {% The two main formation scenarios of binary systems with misaligned discs are the gravitational capture of a passing object in a dense environment, and the fragmentation of the collapsing molecular cloud. Given the low-density environment of the Taurus-Aurigae star-forming region, the first scenario is unlikely for Haro~6-10. The binary system most probably formed via fragmentation of two different parts of the collapsing molecular cloud combined with other dynamical processes related to the cloud and/or the protostars. This can be the explanation also for other binary systems with an infrared companion. } ",
        "introduction": "The formation of binary and multiple systems seems to be a common phenomenon in star formation \\citep[e.g.,][]{Mathieuetal2000}. Surveys of young star-forming regions with high spatial resolution allow a statistical description of the binary parameters and a comparison between regions of various ages and with different environments \\citep[e.g.,][]{Leinertetal1993, Ducheneetal2007}. Binarity has been studied also in young stellar clusters of various ages to investigate its relation to the evolution of protoplanetary discs \\citep[e.g.,][]{ClarkePringle1991a, Bouwmanetal2006}.\\\\ Different scenarios have been proposed to explain the formation of binary systems. The standard scenario supports the hypothesis that a binary or multiple system forms when the core of a molecular cloud fragments during its gravitational collapse. The fragmentation process can be divided into two main classes: direct  \\citep[e.g.,][]{BossBodenheimer1979, BateBurkert1997} and rotational fragmentation \\citep[e.g.,][]{Bonnell1994, BonnellBate1994a, BonnellBate1994b, Burkertetal1997}. While the direct fragmentation strongly depends on the initial density distribution in the molecular cloud, this is not the case for rotational fragmentation, which is caused by asymmetric instabilities in a rotating disc or ring. If the proto-binary accretes (only) gas with low angular momentum, a disc is formed only around the primary star while the secondary may have none. If instead the secondary developed its own disc, also the primary will have a disc \\citep{BateBonnell1997}.  According to the two different scenarios of binary formation, the resulting circumstellar discs may have different inclinations. When the binary is formed via rotational fragmentation the discs are expected to be preferentially aligned. If, however, the stellar components and their own discs form directly from a collapsing core \\citep[`prompt initial   fragmentation';][]{Pringle1989}, this can lead to two misaligned discs \\citep{Bateetal2000}. Under more complicated initial conditions even a small cluster can be formed from such a core. Moreover, interactions in these clusters between stars with discs can cause a dissipative encounter \\citep[e.g.,][]{ClarkePringle1991a}. This may lead to the formation of binary or multiple systems via capture and thus to binaries with misaligned discs. In dense star-forming regions, a massive accretion disc may even cause the capture of a passing star \\citep[][]{ClarkePringle1991b}. On the other hand, strong tidal interactions can lead to aligned discs.  Another possibility to have misaligned discs is that the gravitational interaction of a passing object causes the disc's tilt (and hence also the direction of a possible jet) in a binary system \\citep[][]{Bateetal2000}. \\cite{Heller1993} estimated that in one million years, in a small cluster, a few percent of the stellar systems can experience a tilt of a few degrees during the star encounters. The effect of stellar encounters in different cluster environments has been investigated by \\citet[][]{Olczaketal2010}. They found that the encounter rate  remains unaffected by the size of the stellar population, while it strongly depends on the cluster density: in clusters less dense and less massive than the ONC, massive stars dominate the encounter-induced disc-mass loss, whereas in denser and more massive clusters the disc-mass removal is controlled by low-mass stars.  In systems near the end of their accretion phase, the infall of a small amount of material with a different angular momentum to the orbit can cause a misalignment of discs originally aligned \\citep[e.g.,][]{Bateetal2000}.  Therefore, from the alignment of circumstellar discs it is possible to understand the dynamics that most likely took place during the star-formation process.  Studying disc alignment requires observations of protoplanetary discs around young binary stars with high spatial resolution. Previous studies determined the inclination by means of polarimetry \\citep[e.g.,][]{Moninetal1998, Wolfetal2001, Moninetal2006}, orientation of the jets \\citep[e.g.,][]{Eisloffeletal1996, Davisetal1997}, and direct observations of the discs \\citep[e.g.,][]{Koresko1998, Stapelfeldtetal1998a}. Another technique, that allows to directly measure size, inclination and orientation of the inner parts of protoplanetary discs is long-baseline interferometry at infrared wavelengths. A recent interferometric study of T Tau showed that the three discs in this triple system are misaligned \\citep{Ratzkaetal2009}. Therefore, the formation process of T Tau probably has been highly dynamic.  We present in this paper a multi-wavelength, high-resolution observational survey of the young binary system Haro 6-10 (GV~Tau, IRAS 04263+2426). Haro 6-10 is a T~Tauri binary system residing within the L-1524 molecular cloud. It is composed of an optically visible source, Haro 6-10~S, and an infrared companion (IRC), Haro~6-10~N. With an angular separation of 1.\\arcsec2, Haro 6-10 N is actually located $\\sim$160~AU north of Haro 6-10~S at the distance of the Taurus-Auruga star-forming region. The binary system was firstly resolved by \\citet{LeinertHaas1989} using speckle interferometry. Haro~6-10~N has a very red spectral energy distribution and is brighter than Haro~6-10~S at wavelengths longer than 4~$\\mu$m. The binary system is variable, particularly in the near-IR, on a timescale as short as a month \\citep{Leinertetal2001}. \\citet{Leinertetal2001} suggest that the variability of Haro~6-10~S is associated to inhomogeneities in the circumstellar material, while the variable behaviour of the northern component is likely associated to a variable accretion mechanism. The first resolved images in the near-infrared of this system, obtained by \\citet{Menardetal1993}, suggest the presence of a flat circumbinary envelope, or disc.  The infrared spectrum of Haro 6-10 between 8 and 13~$\\mu$m has been interpreted by \\cite{Hanneretal1998} by an optically thick plus an optically thin emission, where the optically thin emission could be associated to an envelope, while the optically thick emission is associated to a disc, to the star itself, or to a featureless dust continuum. The Spitzer/IRS spectrum \\citep[between 5-36~$\\mu$m;][]{Furlanetal2008} shows a broad silicate absorption feature at 10~$\\mu$m, but the spatial resolution of the instrument is too low to resolve the binary.  \\citet{Reipurthetal2004} obtained high-resolution 3.6 cm observations with the VLA. In the 3.6\\,cm map the two components of the binary are spatially resolved. At this wavelength the radio continuum emission originates from shock-ionised gas very close to the outflow exciting source \\citep[e.g.,][]{Angladaetal1998}. \\citet{Reipurthetal2004} suggest that Haro~6-10~S is the driving source of the large outflow activity seen in the near-infrared. HCN and C$_2$H$_2$ in absorption towards Haro~6-10~N were firstly detected by \\citet{Gibbetal2007} in the near-infrared using NIRSPEC at the {\\em Keck Telescope} and one year later by \\citet{Doppmannetal2008}, using the same instrument.  In this paper, we are using a multi-wavelength approach, including different techniques, to characterise the binary system Haro 6-10. This paper is organised as follows: in Section~\\ref{obs} we present the archival optical and near-infrared photometric observations, and our mid-infrared spectroscopic and interferometric observations. The data reduction and the results are described in Section~\\ref{datared}. The analysis and the discussion of the results are presented in Sections~\\ref{analysis} and \\ref{discussion}, respectively. Conclusions are drawn in Section~\\ref{conclusions}.   \\begin{table} \\caption{Log of TIMMI~2 observations of Haro~6-10 and of the spectro-photometric standard stars.} \\label{timmi2} \\centering \\begin{tabular}{c c c c} \\hline\\hline \\noalign{\\smallskip} Date       & UT    & Objects  & Airmass\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 2002 Feb.~06 & 00:10 & HD32887  & 1.02\\\\ & 00:41 & Haro~6-10    & 1.71\\\\ & 00:53 & Haro~6-10    & 1.73\\\\ & 01:41 & HD29139  & 1.56\\\\ & 07:46 & HD123139 & 1.11\\\\ \\noalign{\\smallskip}  2002 Dec.~26& 00:42 & HD29139  & 1.82\\\\ & 01:35 & Haro~6-10    & 1.84\\\\ & 01:51 & Haro~6-10    & 1.78\\\\ & 05:00 & HD32887  & 1.07\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table}  \\begin{table*} \\caption{Log of MIDI observations of Haro 6-10~N and S and calibrators observed. $^\\star$: calibrators used to obtain the final calibrated visibilities and fluxes;  $^{\\star\\star}$: quality not good enough to be analysed}. \\label{obs_gvtau} \\centering \\begin{tabular}{c c  r  c  c  c  c} \\hline\\hline \\noalign{\\smallskip} &&&Airmass&Proj.~Baseline&Pos.~Angle&Notes\\\\ Date       &UT & Objects &&[m]&[deg] &\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} &  &        &  &[UT2-UT3]&&\\\\ \\noalign{\\smallskip} 2004 Dec. 28 &00:18&  HD37160      &2.08& 27.61 & 14.61  & Calibrator$^\\star$\\\\%\\noalign{\\smallskip} &01:01&  Haro~6-10~N  & 1.73 &  28.43 & 48.67 & Science\\\\%\\noalign{\\smallskip} &02:38&  Haro~6-10~S  &1.60& 38.13 & 52.15 & Science\\\\%\\noalign{\\smallskip} &03:10&  HD37160      &1.22& 39.98 & 43.62 & Calibrator$^\\star$\\\\%\\noalign{\\smallskip} &04:08&  HD37160      &1.21& 43.61 & 46.24 & Calibrator$^\\star$\\\\%\\noalign{\\smallskip} &05:26&  HD50778      &1.03& 46.43 & 42.72 & Calibrator\\\\%\\noalign{\\smallskip} &06:10&  HD94510      &1.43& 45.17 & 15.53 & Calibrator\\\\%\\noalign{\\smallskip} &07:02&  HD49161      &1.40& 46.61 & 45.68 & Calibrator\\\\%\\noalign{\\smallskip} &07:45&  HD94510      &1.25& 44.00 &32.44 & Calibrator\\\\ \\noalign{\\smallskip} \\noalign{\\smallskip} &   &       &  &[UT3-UT4]&&\\\\ \\noalign{\\smallskip} 2005 Nov. 23 &04:17&  Haro~6-10~N$^{\\star\\star}$  &1.56& 62.13 & 106 & Science\\\\%\\noalign{\\smallskip} &05:22&  Haro~6-10~S  &1.54&  58.60 & 95.6  & Science\\\\%\\noalign{\\smallskip} &05:41&  Haro~6-10~N  &1.56&  56.80 & 92.6  & Science\\\\%\\noalign{\\smallskip} &06:09&  HD25604      &1.61& 50.53 & 85.4 & Calibrator$^\\star$\\\\%\\noalign{\\smallskip} &07:49&  HD33042      &1.23& 61.60 & 116.5 & Calibrator\\\\ \\noalign{\\smallskip}  \\hline \\end{tabular} \\end{table*} ",
        "conclusions": " \\begin{enumerate} \\item Both components of the binary system Haro~6-10 are embedded in a common envelope. The envelope has a dust composition  similar to the interstellar medium. \\item Each component of the system has a circumstellar disc-like structure typical of young stars. \\item The discs are highly misaligned: the northern disc is seen almost edge-on with respect to the line-of-sight and the southern component is surrounded by an almost face-on disc. \\item The most probable scenario is the formation via fragmentation of two different parts of the collapsing molecular cloud combined with other dynamical processes related to the cloud or the protostars. \\end{enumerate}"
    },
    "1107/1107.5315_arXiv.txt": {
        "abstract": "Detailed chemical abundances for twenty one elements are presented for four red giants in the anomalous outer halo globular cluster Palomar 1 ($R_{\\rm{GC}} = 17.2$ kpc, $Z=3.6$ kpc) using high-resolution ($R=36000$) spectra from the High Dispersion Spectrograph (HDS) on the Subaru Telescope. Pal 1 has long been considered unusual because of its low surface brightness, sparse red giant branch, young age, and its possible association with two extragalactic streams of stars---this paper shows that its chemistry further confirms its unusual nature.  The mean metallicity of the four stars, $[\\rm{Fe/H}] = -0.60 \\pm 0.01$, is high for a globular cluster so far from the Galactic center, but is low for a typical open cluster.  The [$\\alpha$/Fe] ratios, though in agreement with the Galactic stars within the $1\\sigma$ errors, agree best with the lower values in dwarf galaxies.  No signs of the Na/O anticorrelation are detected in Pal 1, though Na appears to be marginally high in all four stars.  Pal 1's neutron capture elements are also unusual: its high [Ba/Y] ratio agrees best with dwarf galaxies, implying an excess of second-peak over first-peak s-process elements, while its [Eu/$\\alpha$] and [Ba/Eu] ratios show that Pal 1's contributions from the r-process must have differed in some way from normal Galactic stars.   Therefore, Pal 1 is chemically unusual, as well in its other properties.  Pal 1 shares some of its unusual abundance characteristics with the young clusters associated with the Sagittarius dwarf galaxy remnant and the intermediate-age LMC clusters, and could be chemically associated with the Canis Majoris overdensity; however it does not seem to be similar to the Monoceros/Galactic Anticenter Stellar Stream. ",
        "introduction": "\\label{sec:Intro} Ever since its discovery by \\citet{Abell}, Palomar 1 (Pal 1) has been tentatively classified as a globular cluster (GC), primarily because of its location high above the Galactic plane ($Z=3.6$ kpc; \\citealt{Harris}, 2010 edition).  However, Pal 1 has several anomalous characteristics that have cast doubt upon this classification. For example, Pal 1 appears to be unusually young for an outer halo GC. Using standard isochrone fits to Pal 1's ground-based color-magnitude diagram (CMD), \\citet{RosenbergA} estimated an age between $6.3$ and $8$ Gyr, while \\citet{Sarajedini} found an age between 4 and 6 Gyr with HST photometry.  This F606W ($\\sim$V), F814W ($\\sim$I) HST CMD of the central field of Pal 1 from \\citealt{Sarajedini} (see Figure \\ref{fig:CMD}) shows the main sequence turn-off around F606W $\\approx 19.4$ ($V\\approx 19.6$), and a possible detection of a red horizontal branch near F606W $\\approx 16.4$ ($V\\approx 16.6$).  Thus, $\\Delta V^{\\rm{HB}}_{\\rm{MSTO}} \\approx 3$, making the cluster younger than about 7 Gyr (with the method of \\citealt{Chaboyer}).  These estimates make Pal 1 one of the youngest GCs, along with Ter 7, Pal 12, and Whiting 1, whose ages range from $7-9$ Gyr (e.g. \\citealt{Buonanno1998}; \\citealt{SalarisWeiss}; \\citealt{Carraro}; \\citealt{Dotter2010}).  \\begin{figure} \\epsscale{1.2} \\plotone{Pal1CMDandTargets.eps} \\caption{An HST color-magnitude diagram (roughly V, V-I) for the central field of Pal 1 from \\citet{Sarajedini}.  The observed targets are circled.  Also shown are fitted isochrones from the Dartmouth Stellar Evolution Database \\citep{DSED} for ages of 4, 5, and 6 Gyr. Values of $[\\rm{Fe/H}] = -0.6$, $[\\alpha/\\rm{Fe}]=0.0$, $(m-M)_{V} = 16.27$, and $E(B-V) = 0.20$ are adopted, as discussed in the text\\label{fig:CMD}} \\end{figure}  Pal 1 also lacks luminous, evolved stars.  The CMD in Figure \\ref{fig:CMD} shows a sparsely populated red giant branch and a barely detectable horizontal branch.  Furthermore, near-infrared medium-resolution spectroscopy of the \\ion{Ca}{2} triplet (CaT) lines in four of these red giants has shown that Pal 1 is fairly metal-rich for an outer halo GC, with an average $[\\rm{Fe/H}] = -0.6$ (on the Zinn-West scale, using the standard Galactic CaT line strength to metallicity conversion; \\citealt{RosenbergB}).  The main sequence is also poorly populated, and therefore the entire cluster has an extremely low total luminosity for a GC ($M_V = -2.52$; \\citealt{Harris}, 2010 edition).  Clearly Pal 1 does not conform to the characteristics of traditional Galactic GCs, suggesting either that Pal 1 is not a GC or that it is not Galactic.  If Pal 1 has been misclassified as a GC, another possibility is that it could be an open cluster (OC)---this option seems unlikely due to the combination of its position above the Galactic plane, its high concentration parameter ($c=2.57$; \\citealt{Harris}, 2010 edition), and its age (although Pal 1's age \\textit{is} comparable to that of the oldest known OCs, e.g. \\citealt{VandenbergStetson}).  However, Pal 1's brightness and age place it in greater agreement with the OCs than with the GCs (see Figure 3 in \\citealt{Carretta2010}).  The ambiguity surrounding the OC vs. GC classification leads to a need for a more rigorous distinction between the two.  From observations of many Galactic GCs, \\citet{Carretta2010} defined a GC as a cluster that shows signs of a second generation of stars through a Na/O anticorrelation.  Under this definition, if Pal 1 does \\textit{not} show signs of a Na/O anticorrelation, then it cannot be considered a classical GC.\\footnote{In the past, Pal 1's unusual characteristics have led to a designation as a \\textit{transitional} cluster, i.e. an object that resembles something between classical globular and open clusters, e.g., \\citet{Ortolani}.}  Alternatively, Pal 1 could have originated in a dwarf galaxy and been accreted by the Milky Way during a merger. The latter situation is plausible, as several other Galactic GCs, notably the young clusters listed above, have been linked to dwarf galaxies.  M54, Arp 2, Ter 7, and Ter 8 lie on a stream associated with the Sagittarius dwarf spheroidal (Sgr dSph) galaxy \\citep{DaCostaArmandroff}; Pal 12 has been kinematically linked to the Sgr dSph through proper motion studies \\citep{Dinescu}; Whiting 1 is undergoing tidal stripping and appears to lie in a stream of M giants from the Sgr dSph \\citep{Carraro}; and Rup 106 has been tentatively associated with the Magellanic Clouds because of its location along the Magellanic stream \\citep{LinRicher}.  Pal 1 does appear to lie along several stellar streams that may be extragalactic but have not yet been associated with any galaxies.  One such stream, known both as the Monoceros stream and as the Galactic Anticenter Stellar Stream (GASS), appears to be from a disrupted dwarf satellite (e.g. \\citealt{Sollima}); Pal 1's position, velocity, and metallicity make it a possible member of this stream \\citep{Crane}.  The Canis Major overdensity, which seems to be distinct from the GASS \\citep{Chou}, has also been associated with Pal 1, again because of the cluster's coincidental location \\citep{Martin}. Finally, Pal 1 has also been tentatively linked to the ``Orphan Stream'' \\citep{Belokurov}, though its radial velocity and metallicity make it an unlikely member \\citep{Newberg}. Though no obvious association with a particular stream has been confirmed, tidal tails around Pal 1 have recently been discovered, suggesting that the cluster itself is being dissipated in the tidal field of the Galaxy \\citep{NiedersteOstholt}.  In the absence of an obvious host galaxy, detailed chemical abundances can be used to examine if Pal 1 has an extragalactic origin.  It has been well established (e.g. \\citealt{Tolstoy}, \\citealt{Venn2004}) that metal-rich stars in Local Group dwarf galaxies have distinctly different chemical abundances from stars in the Milky Way, at a given metallicity (the differences are not so obvious by $[\\rm{Fe/H}] \\sim -2$).  These chemical variations between the environments are due to different star formation histories and efficiencies, which are linked to the mass of the galaxy and the properties of its environment. Therefore, the differences are not just between dwarf galaxies and normal galaxies---in fact, dwarf galaxies have unique chemical evolutions, meaning that it is possible to chemically link a star to its host galaxy (e.g. \\citealt{Tolstoy}, \\citealt{Freeman}). These chemical differences are not just seen in field stars, they are also seen in GCs; for example, \\citet{Pritzl} showed that most Galactic GCs have similar chemical abundances to Galactic field stars at the same metallicity, with Pal 12, Ter 7, and Rup 106 as notable exceptions.  Detailed, high-resolution chemical abundances have been used to trace the origin of several globular clusters.  In particular Pal 12 \\citep{Cohen} and Ter 7 (\\citealt{Tautvaisiene}, \\citealt{Sbordone}) have [$\\alpha$/Fe] ratios that follow the Sgr dSph field star abundance trends and are clearly separated from Galactic field stars, supporting the hypothesis that they were captured from the Sgr dSph. The more metal-poor Sgr clusters Arp 2 and Ter 8, however, are not chemically distinct from the Galaxy \\citep{Mottini}, which shows that the [$\\alpha$/Fe] deviation can be metallicity- (and presumably age-) dependent.  If Pal 1 has an extragalactic origin, then its metallicity of $[\\rm{Fe/H}]\\approx -0.6$ puts it in a regime where [$\\alpha$/Fe] and other chemical abundance ratios should be useful in determining its origin.  Here we present a detailed analysis of high-resolution spectra for four red giant stars in Pal 1, taken with the Subaru Telescope. \\citet{MonacoPal1} have completed a chemical analysis of one star in Pal 1 (Pal 1-I), also with spectra taken at Subaru, but with independent observations; they find that their single star shows similar abundance patterns to Galactic open clusters. \\citet{Saviane} further show that its chemistry implies that Pal 1 could be associated with the Canis Major overdensity.  With four stars, our analysis allows us to determine the chemical abundance ratios with higher precision, and investigate any star-to-star variations in the abundances that could be linked with chemical or cluster evolution effects. Section \\ref{sec:Observations} outlines the observations and the data reduction, while Section \\ref{sec:EWs} examines our method for measuring equivalent widths.  In Section \\ref{sec:ModelAtms} we discuss the model atmospheres, particularly how the atmospheric parameters are derived. Section \\ref{sec:Abundances} presents the methods for element abundance determinations, while Section \\ref{sec:Results} examines the significance of these results with respect to Galactic stars and clusters. Finally, Section \\ref{sec:Discussion} examines the origins of Pal 1. ",
        "conclusions": "\\label{sec:Conclusion} Detailed chemical abundances have been found for twenty one elements in four red giant stars in the unusual GC Palomar 1.  Our findings are summarized as follows:  \\begin{itemize} \\item Pal 1 is an outer halo, young ($5\\pm 1$ Gyr) cluster whose age, metallicity, location, and structural parameters distinguish it from the standard Galactic globular or open clusters.  However, Pal 1 does appear to be similar to the known extragalactic GCs and OCs that have been accreted during mergers (i.e. from the Sgr dSph).  \\item Pal 1's young age and single stellar population suggest a resemblance to the low-mass intermediate-age LMC clusters.  Chemical abundances further suggest that Pal 1 may have shared a similar formation history.  \\item Pal 1 shows several unusual chemical characteristics, including \\begin{itemize} \\item Lower [$\\alpha$/Fe] ratios than Galactic stars of the same [Fe/H] \\item A lack of evidence for a Na/O anticorrelation, though marginally high Na is found \\item Similar [Fe-peak/Fe] ratios to Galactic stars \\item An excess of second-peak s-process neutron capture elements over first-peak elements \\item Low [Ba/Eu] and high [Eu/$\\alpha$] values that suggest Pal 1's host galaxy had an additional r-process site (possibilities could be inhomogeneous mixing, an effectively truncated IMF, i.e. one that is missing the most massive stars, or variable nucleosynthetic yields). \\end{itemize}  \\item Chemically, Pal 1 does not behave like the typical Galactic bulge/disk GCs or the old, metal-poor OCs.  The closest agreement seems to be with the Sgr clusters Pal 12 and Ter 7 and the LMC intermediate-age clusters; however the Na, Al, Cu, and Zn abundances do not agree with the LMC or Sgr clusters.  \\item Comparing the [X/Fe] of the Pal 1 stars to those of stars in known streams show that \\begin{itemize} \\item It is unlikely that Pal 1 originated in the Galactic Anticenter Stellar Stream, given the differing [La/Fe] ratios \\item Pal 1 may have originated in the Canis Major overdensity, if the [X/Fe] ratios of the three possible CMa stars are extrapolated to slightly more metal-poor stars. \\end{itemize}  \\end{itemize}  Overall, we conclude that Pal 1 likely had an extragalactic origin, though its chemistry remains unique compared to the known globular and open clusters."
    },
    "1107/1107.0195_arXiv.txt": {
        "abstract": "We present a study of the core of the \\fekalfa emission line at $\\sim 6.4$~keV { in a sample of type II Seyfert galaxies observed by the \\chandra High Energy Grating (HEG). The sample consists of 29 observations of 10 unique sources. We present measurements of the \\fekalfa line parameters with the highest spectral resolution currently available. In particular, we derive the most robust intrinsic line widths for some of the sources in the sample to date. We obtained a weighted mean FWHM of $2000 \\pm 160 \\ \\rm km \\ s^{-1}$ for 8 out of 10 sources (the remaining sources had insufficient signal-to-noise). From a comparison with the optical emission-line widths { obtained from spectropolarimetric observations}, we found that the location of \\fekalfa line-emitting material is a factor of $\\sim0.7-11$ times the size of the optical BLR. Furthermore, compared to 13 type I AGNs for which the best \\fekalfa line FWHM constraints were obtained, we found no difference in the FWHM distribution or the mean FWHM, and this conclusion is independent of the central black hole mass. This result suggests that the bulk of the \\fekalfa line emission may originate from a universal region at the same radius with respect to the gravitational radius, $\\sim3\\times10^4$ $r_g$ on average. By examining the correlation between the \\fekalfa luminosity and the [O IV] line luminosity, { we found a marginal difference in the \\fekalfa line flux between type I and type II AGNs, but the spread in the ratio of $L_{\\rm Fe}$ to \\loiv~ is about two orders of magnitude. Our results confirm the theoretical expectation that the \\fekalfa emission-line luminosity cannot trivially be used as a proxy of the intrinsic AGN luminosity, unless a detailed comparison of the data with proper models is applied.}} ",
        "introduction": "\\label{intro} Both type I and type II active galactic nuclei (AGN) are known to exhibit a narrow (FWHM~$<10,000 \\ \\rm km \\ s^{-1}$) \\fekalfa fluorescent emission line at $\\sim 6.4$~keV in their X-ray spectrum (e.g. Sulentic \\etal 1998; Lubi\\'{n}ski \\& Zdziarski 2001; Weaver, Gelbord, \\& Yaqoob 2001; Perola \\etal 2002; Yaqoob \\& Padmanabhan 2004 (hereafter YP04); Levenson \\etal 2002, 2006; { Jiang et al. 2006;} Winter \\etal 2009; Shu, Yaqoob \\& Wang 2010, hereafter paper I). The line profile of the \\fekalfa core is important for probing its origin and it can provide unique information on the dynamics and physical state of the line-emitting region (Yaqoob et al. 2001, 2003, 2007). While it is widely accepted that the narrow \\fekalfa line cores are produced in cold, neutral matter far from the nucleus, the exact location and distribution of the line-emitting gas still remain uncertain (see Paper I, and references therein). Nandra (2006) first examined the relation between the \\fekalfa and optical H$\\beta$ line widths, which can potentially give a direct indication of the location of the \\fekalfa line-emitting region relative to the optical broad-line region (BLR). However, the results were ambiguous, and the data allow for an origin of the \\fekalfa line anywhere from the outer regions of an accretion disk, the BLR, and a parsec-scale torus. Meanwhile, Bianchi et al. (2008) reported a meaningful comparison between the \\fekalfa and H$\\beta$ line widths in NGC 7213, and found the FWHM of both lines are consistent with each other ($\\sim$2500 km s$^{-1}$), implying a BLR origin of the \\fekalfa emission line. Using the high-energy grating (HEG) on the \\chandra \\hetg ({\\it High Energy Transmission Grating Spectrometer}, see Markert et al. 1995), which affords the best spectral resolution currently available in the Fe K band (at 6.4 keV is $\\sim$39 eV, or $\\sim$1860 km s$^{-1}$ FWHM), in Paper I we presented a more thorough and comprehensive study of the \\fekalfa line core emission in a large sample of type I Seyfert galaxies (see also YP04). We measured the intrinsic width of the narrow \\fekalfa line core and obtained a weighted mean of FWHM = 2060 $\\pm$ 230 km s$^{-1}$. A comparison with the optical emission-line widths suggested that there may not be a universal location of the \\fekalfa line-emitting region relative to the BLR, consistent with the results of Nandra (2006).  Our results in fact showed that the location of the \\fekalfa line emitter relative the BLR appears to be different from source to source.   However, one must be cautious in explaining the origin of the narrow \\fekalfa core, especially for measurements made with instruments that have lower spectral resolution than the \\chandra \\hetg (e.g., see Liu et al. 2010 for NGC 5548). The uncertainty in the intrinsic line width measurements is usually large, and part of the narrow component may have a contribution from an underlying broad line in some (if not all) unobscured AGNs (e.g., Lee et al. 2002; YP04; Nandra 2006). However, in cases when the narrow \\fekalfa core is unresolved even with the \\chandra HEG, such contamination is not an issue, and from the upper limits on the line widths we can place strong constrains on the origin of the intrinsically narrow \\fekalfa line core (see Paper I). In the paradigm of the unification model (Antonucci 1993), the narrow \\fekalfa line emission in type II AGNs is expected to be produced in the same material as in type I nuclei. Therefore, one of the things that we would like to know is whether there is any systematic difference in the origin of the \\fekalfa line in type I and type II AGNs. In addition, for a given geometry and model, one would expect a particular relation between the equivalent width (EW) of the \\fekalfa line, $N_H$, and the orientation of line-emitting structure (e.g., Murphy \\& Yaqoob 2009; Ikeda, Teramshima, \\& Awaki 2009; Yaqoob et al. 2010; Brightman \\& Nandra 2011). Thus a comparison of \\fekalfa emission-line luminosities and EWs in type I AGNs versus type II AGNs can potentially offer basic tests of some of the predictions of AGN unification scenarios.  In this paper, we present a study of the narrow \\fekalfa emission line in a sample of type II AGNs observed with the Chandra \\hetgp. The observations and spectral fitting are described in Section 2. In Section 3, we present the results for the properties of the core of the narrow \\fekalfa emission line and their implications. In Section 4, we summarize our results and findings. Throughout this paper, we adopt a cosmology of $\\Omega_M=0.3, \\Omega_{\\lambda}=0.7$, and H$_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\subsection{Spectroscopy in the Fe K Band} \\label{hegspec} Figure 1 shows the HEG spectrum for each source in the Fe K region, corrected for instrumental efficiency and cosmological redshift. The spectra are binned at 0.01 \\AA , similar to the HEG FWHM spectral resolution of 0.012 \\AA. The statistical uncertainties correspond to 68\\% confidence Poisson errors, which we calculated using equations (7) and (14) in Geherls (1986) that approximate the upper and lower errors respectively. For sources which were observed more than once, the time-averaged spectra are shown. It can be seen that in some spectra, emission lines from different ionized species of Fe are evident, in addition to the narrow \\fekalfa line at $\\sim$6.4 keV, but the latter is always the strongest line. Although in the present paper we are concerned only with the \\fekalfa line core centered at $\\sim$ 6.4 keV, overlaid on the spectra in Fig. 1 are vertical dashed lines marking the expected positions of the \\fexxv He-like triplet lines (the two intercombination lines are shown separately), \\fexxvip~Ly$\\alpha$, Fe I K$\\beta$, and the neutral Fe K-shell threshold absorption-edge energy. The values adopted for these energies were from NIST2 (He-like triplet); Pike et al. 1996 (\\fexxvip~ Ly$\\alpha$); Palmeri et al. 2003 (Fe I K$\\beta$ ), and Verner et al. 1996 (Fe K edge). We emphasize that the emission from higher ionization states of Fe has little effect on the measurements of the intrinsic width of the \\fekalfa line. For example, for NGC 1068 the spectral plot in the Fe K band shows a broad emission feature on the blue side of the \\fekalfa line peak, which is probably due to multiple species of ionic Fe. {The occurrence of highly ionized Fe emission lines in AGN X-ray spectra has been noted since the $ASCA$ observations, and they are extensively studied recently with \\xmm and $Suzaku$ (e.g., Bianchi et al. 2005; Fukazawa et al. 2011). } However, we can confirm from the 99\\% confidence contour of \\fekalfa line intensity versus FWHM (Figure 2), that the \\fekalfa line width is still well-constrained (less than $\\sim$3500 km s$^{-1}$). This conclusion is also consistent with the physical width of the narrow core of the \\fekalfa line in the spectral plot that is apparent simply from visual inspection.   The best-fitting \\fekalfa emission-line parameters for each spectrum are shown in Table 2. We do not give the best-fitting values of $\\Gamma$ or $N_{H}$ in Table 2 because the values derived using the simplistic continuum model are not physically meaningful but are simply parameterizations. We obtained a weighted mean line-center energy of 6.397$\\pm$0.001 keV for the 24 observations of 8 sources. At 99\\% confidence (for three free parameters), none of the AGNs has a line centroid energy greater than $\\sim$6.41 keV, indicating that the line-emitting matter is cold and essentially neutral, consistent with results presented in Paper I for a sample of type I AGNs. Here, and hereafter, for the calculation of the weighted mean of any quantity with asymmetric errors, we simply assumed symmetric errors, using the largest 68\\% confidence error from spectral fitting.   \\subsection{The Intrinsic Width of the \\fekalfa Line Emission}  From our spectral fits (Table 2) and the joint confidence contours of \\fekalfa line intensity versus FWHM (Figure 2 and 3), we found that at 99\\% confidence, the \\chandra HEG resolves the narrow component of the \\fekalfa emission in 8 out of 10 sources. The weighted mean FWHM of the \\fekalfa line cores is $2000 \\pm 160 \\ \\rm km \\ s^{-1}$. { The fact that this mean is approximately equal to the HEG FWHM spectral resolution is not indicative of a calibration bias. This is supported by the fact that the FWHM of the Fe~K$\\alpha$ line in Circinus is {\\it less} than the HEG FWHM at a confidence level of greater than 99\\%. In addition, the spectral resolution for gratings does not degrade with time because it is determined principally by spatial dispersion, and the resolution is well-established from bright Galactic sources with narrow lines (http://space.mit.edu/CXC/calib/hetgcal.html). Note that the weighted mean value is consistent with the straight mean of the FWHM, of $2510 \\pm 160 \\ \\rm km \\ s^{-1}$. } The measurements of the intrinsic width of the line can potentially be used to constrain the location of the medium responsible for the core of the \\fekalfa emission line (e.g., Yaqoob et al. 2001). Comparing the \\fekalfa line FWHM with that of the optical broad emission lines (e.g., H$\\alpha$ and/or \\hbeta lines) can give a direct indication of the location of the \\fekalfa line-emitting region relative to the optical broad-line region (Nandra 2006, Bianchi et al. 2008, Paper I). To place limits on the location of the \\fekalfa emitter relative to the BLR, we compiled from the literature the width of optical broad emission line from spectro-polarimetric observations for five sources in our sample, and the values of the H$\\alpha$ or \\hbeta FWHM are listed in Table 2. { The near-infrared broad lines have not been taken into account in the present paper, as they are possibly produced in a region that differs from the BLR in density, and/or in the amount of extinction (e.g., Ramos Almeida \\etal 2008; Landt \\etal 2008). However, in cases where there is no spectropolarimetric data at all, the FWHM of near-infrared lines might still yield some useful information}\\footnote{ We searched in the literature and found that only one source, NGC 6240, has NIR broad-line Br$\\alpha$ measurement reported, with FWHM=1800$\\pm$200 km/s (Cai et al. 2010). This value is consistent with the width of the Fe K line at 99\\% confidence.}.  Figure 3 shows the 99\\% confidence contours of the \\fekalfa line intensity versus the {\\it ratio} of the \\fekalfa FWHM to optical line FWHM. We found that at the two-parameter 99\\% confidence level, the FWHM ratio lies in the range of $\\sim0.3-1.2$. At 99\\% confidence, the \\fekalfa line widths in 3 out of 5 sources are less than that of optical lines, suggesting that the observed \\fekalfa emission is not likely produced in a BLR where the optical lines are produced. The remaining two sources ({ NGC 4388 and NGC 1068}) could have consistent width for both lines. This would be expected in at least one source, NGC 4388, which is known to have significant column density variations (Elvis et al. 2004), if its BLR acts as absorber and emitter seen in the X-ray band. In fact, there is now significant evidence that the BLR may act as an X-ray absorber, like in the cases of NGC 1365 and NGC 7582 (e.g., Risaliti et al. 2005; Bianchi et al. 2009). { However, if the BLR is obscured by a fully covering column density greater than $\\sim 5 \\times 10^{23} \\ \\rm cm^{-2}$ ({ as in the cases of NGC 4507, Mrk 3 and the Circinus} ) any Fe~K$\\alpha$ line component from the BLR will be greatly diminished, or undetectable. If the BLR is obscured by such a large column density, the only way it would be possible to observe an Fe~K$\\alpha$ line component from the BLR is if the BLR is not fully covered.} { The fact that the \\fekalfa line width of Compton-thick source NGC 1068 is consistent with that of the polarized \\hbeta line (within 99\\% errors) could be simply attributed to the large error in the \\fekalfa line width. Alternatively, this could indicate that the heavy obscuration which blocks the BLR has comparable size with that of the BLR. }%   Since Keplerian velocities are inversely proportional to the square root of the orbit size, { the observed distribution of the ratio of the \\fekalfa to optical line FWHM} implies that the \\fekalfa line-emitting region size could be a factor $\\sim 0.7-11$ times larger than the optical line-emitting region. We obtained a weighted mean of 0.57$\\pm$0.05 for the ratios of the \\fekalfa line FWHM to the optical line FWHM, corresponding to the \\fekalfa line-emitting region being, on average, $\\sim$3 times the size of BLR. However, for AGNs in general, there may be no universal location of the \\fekalfa line-emitting region within a factor of $\\sim$1--10 times size of BLR. However, as we will show later, the \\fekalfa line emitter may be associated with a universal location in terms of gravitational radius ($r_g$, where $r_g=GM_{\\rm BH}/c^2$). If the \\fekalfa line emission has a significant contribution from the putative obscuring torus that is required by AGN unification models, our results show that the size-scale of the torus may be smaller than traditionally thought. Note that Gaskell, Goosmann, \\& Klimek (2008) argue that there is considerable observational evidence that the BLR itself has a toroidal structure, and that there may be no distinct boundary between the BLR and the classical parsec-scale torus.  One of the interesting things that we can investigate is whether the material responsible for producing the narrow \\fekalfa line in type II AGNs systematically differs from that in type I objects. Of the AGNs in the sample reported in Paper I, we identified { thirteen} type I AGNs (including 3 moderately obscured sources with weak broad Balmer lines, formally classified as intermediate type AGNs) that provided the very best statistical constraints on the \\fekalfa line FWHM (as indicated by the two-parameter confidence contours of line flux versus FWHM; see Figure 6 in Paper I). The weighted mean FWHM of the \\fekalfa core { is 2170$\\pm$220 km $^{-1}$}, in good agreement with the measurement for type II AGNs in the present sample. { However, we found a slightly larger FWHM of the \\fekalfa line (3620$\\pm$220 km $^{-1}$) for type I AGNs if the straight mean is used, but as we will show below, it may be due to a bias towards the measurements with lower quality data.} We show in Figure 4 the distribution of \\fekalfa line FWHM for the current sample (solid line), compared with the { thirteen} type I AGNs in Paper I (dashed line). It can be seen that both histograms are not Gaussian and are not strongly peaked. In addition, we found that the measurement of the FWHM in one AGN (namely MCG -6-30-15) deviated significantly from the distribution of the rest of the sources. From our empirical analysis, we obtained FWHM=11880$^{+4650}_{-4030}$ km s$^{-1}$ (see Paper I) for this object, and the corresponding 90\\% confidence lower limit (for three free parameters) is $\\sim$6000 km s$^{-1}$. Note that MCG -6-30-15 has the strongest and broadest \\fekalfa line yet observed in an AGN. The larger FWHM obtained from our empirical analysis could be attributed to the \\fekalfa core from underlying disk-line component, and/or a complex continuum (e.g., Miller et al. 2008). Note that Young et al. (2005) analyzed the \\chandra HEG spectrum by using more complex models including both disk-line and narrow \\fekalfa core emission, and obtained an FWHM$<$ 4700 km s$^{-1}$ for the narrow component, which is consistent with those from the other sources in our sample. However, as the narrow \\fekalfa line EW is relatively low ($\\sim$60 eV, see Paper I), we cannot tell whether the line profile is affected by the complex continuum, or whether it is intrinsically broader than the other AGNs . Future X-ray missions, such as {\\it Astro--H}, which has much higher spectral resolution, will help to measure the true profile of the narrow \\fekalfa emission line in this object. { Even considering }the one AGN mentioned above, there is not a significant difference in the distribution of FWHM for the two subsamples. A Kolmogorov-Smirnov (K-S) test shows that the probability that the type I ({ including MCG -6-30-15}) and type II AGNs are drawn from the same parent population is 0.83. Therefore, it appears that there is no difference in the origin of the \\fekalfa line in type I and type II AGNs, which is consistent with the predictions of the unification model.  Figure 5 shows the relationship between the \\fekalfa line FWHM and the black hole mass (see Table~2 caption for references). Open and solid circles denote type I and type II AGNs, respectively. We see that all FWHM values, within the statistical errors, are consistent with a constant, independent of the mass of black hole. Assuming that the line originates in material that is in a virialized orbit around the black hole, we can estimate the distance, $r$, of the line-emitting material to the black hole using the relation $GM_{\\rm BH}= r\\langle v^2\\rangle$. Assuming that the velocity dispersion is related to FWHM velocity as $\\langle v^2\\rangle$=$3\\over4$$v^2_{\\rm FWHM}$ (Netzer et al. 1990), we can obtain $r=(4c^{2})/(3v^{2}_{\\rm FWHM})r_g$, where $r_g=GM/c^2$. Using the weighted mean of FWHM $\\sim$2000 km s$^{-1}$, we find that the $r\\sim3\\times10^4r_g$, larger than the typical size of the BLR (e.g., Peterson et al. 2004). Therefore, our results seem to support the bulk of the \\fekalfa line production arising from a region that is appears to be located at a universal distance with respect to the gravitational radius, which is controlled by central black hole mass. Note that Nandra (2006) also examined the relation between the FWHM of \\fekalfa line and the black hole mass, but their results were ambiguous. The reason why they did not find the result that we found is possibly due to the fact that some of the sources in their sample had only poor quality and/or problematic \\fekalfa line-width measurements. Including such sources could have obscured the underlying result that $r/r_{g}$ may be the key factor that determines the location of the line-emitting region. However, it is important to note that even with superb \\chandra HEG spectral resolution, there may be blending from a Compton-shoulder component and/or multiple low ionization states of Fe (e.g., Yaqoob \\& Murphy 2011), so that the true line width may be less than the FWHM deduced from our simplistic model-fitting.  \\subsection{What can we learn from the \\fekalfa Line Emission? } From a theoretical point of view, the flux or luminosity of the \\fekalfa line ($L_{\\rm Fe}$) is nontrivial to calculate, as it depends on a number of factors, including geometry, orientation, covering factor, element abundance, and column density of the line-emitting material. Using Monte Carlo simulations of a toroidal X-ray reprocessor model, Murphy \\& Yaqoob (2009, see also Yaqoob et al. 2010) showed that geometrical and inclination-angle effects become important for $N_H\\gtrsim10^{23}$ cm$^{-2}$ for the observed EW and line flux. For a given covering factor and set of element abundances, a toroidal structure with a column density of  greater than $\\sim 5 \\times 10^{24}$ cm$^{-2}$, observed at an edge-on inclination could produce an \\fekalfa emission line flux that is an order of magnitude or more weaker than a face-on orientation (Yaqoob et al. 2010). In particular, the \\fekalfa line luminosity is not simply a linear function of $N_H$, and it has a maximum value for a column density in the range $\\sim3-8\\times10^{23}$ cm$^{-2}$, depending on the inclination angle. Therefore, the \\fekalfa line luminosity cannot be trivially used to measure the intrinsic AGN luminosity ($L_{\\rm AGN}$). The relation between $L_{\\rm Fe}$ and $L_{\\rm AGN}$ is more complicated than we would expect from optically-thin matter, and it is strongly model dependent, and in particular it depends on covering factor, inclination angle and column density. We emphasize that the column density for producing the \\fekalfa line does not refer to the line-of-sight value, but rather to a value that corresponds to the angle-averaged flux over all incident X-ray continuum radiation (see Murphy \\& Yaqoob 2009).  Recently, Liu \\& Wang (2010) confirmed that $L_{\\rm Fe}$ and $L_{\\rm AGN}$ are not simply related. They compared the narrow \\fekalfa line emission between type I and type II AGNs, using data obtained with \\xmmp, and found that statistically,  { the \\fekalfa line luminosities in Compton-thin and Compton thick type II AGNs are about 2.7 and 5.6 times lower than that in type I sources, respectively. They therefore proposed that different correction factors should be applied if one uses the \\fekalfa line emission to estimate the AGN's intrinsic luminosity.} { To examine the relation between $L_{\\rm Fe}$ and $L_{\\rm AGN}$ for our \\chandra grating sample, we plot in Figure 6 (upper panel) the \\fekalfa line luminosity ($L_{\\rm Fe}$), versus the [O IV] 25.89 $\\mu$m line luminosity (\\loiv, which is claimed to be an intrinsic AGN luminosity indicator, see Mel\\'{e}ndez et al. 2008a). The type II AGNs are shown with solid circles, while open circles represent 24 type I AGNs for which the [O IV] line luminosity is available from literature. Solid stars denote tentative Compton-thick sources, based on previously reported measurements of N$_{\\rm H}$ (e.g., Treister, Urry \\& Virani 2009, and references therein). The right-hand panel shows the distribution of $L_{\\rm Fe}$ for both type I (dotted line) and type II (solid line) AGNs. It can be seen that the distribution of the \\fekalfa line luminosity in type II AGNs is not significantly different from that in type I AGNs. A K-S test shows the probability that both samples were drawn from the same parent population is 0.21.} { However, when normalizing the \\fekalfa line luminosity by the [O IV] line luminosity ($L_{\\rm Fe}$/\\loiv, the lower panel of Figure 6), we find a marginal difference in the distributions of the $L_{\\rm Fe}$/\\loiv~ratio for type I and type II AGNs, but the significance level is not high (at 93\\% confidence), possibly due to the small size of the current \\chandra grating sample. The dotted and dashed horizontal lines show the means of $L_{\\rm Fe}$/\\loiv~ for Liu \\& Wang's type I and type II subsamples, respectively. It seems that the distribution of the $L_{\\rm Fe}$/\\loiv~ratio for our \\chandra grating sample shows a similar pattern as Liu \\& Wang's sample. However, we note that there is significant overlap between the type I and type II AGNs with intermediate values of the $L_{\\rm Fe}$/\\loiv~ratio (see also Liu \\& Wang 2010, Figure 6), suggesting the complex dependencies of the \\fekalfa line emission on the geometry and physical properties of the line-emitting material. In the context of a toroidal geometry (e.g., Murphy \\& Yaqoob 2009), the sources in the overlap region of the $L_{\\rm Fe}$/\\loiv~ plot may however indicate that they constitute a selected distribution of N$_H$ and inclination angle (e.g., Circinus galaxy, Yang et al. 2009), but a larger sample is required to make a detailed comparison. }  Although in our study we cannot determine some of the critical parameters of the \\fekalfa line emitter, such as global covering factor, $N_H$, and the inclination angle of the X-ray reprocessor, we can still make interesting statements from the plot of \\loiv~vs. $L_{\\rm Fe}$/\\loiv. It is apparent that the ratio of $L_{\\rm Fe}$/\\loiv ~appears to have a dispersion of 2 to 3 orders of magnitude or so. Regardless of AGN type, the sources located at the top end of the $L_{\\rm Fe}$/\\loiv~ ratio distribution are expected to be those objects that are observed at a face-on, or near face-on, inclination angle, with moderate $N_H$ ($\\sim0.5-2\\times10^{24}$ cm$^{-2}$). The sources located in the region of the lowest values of $L_{\\rm Fe}$/\\loiv~should correspond to cases with higher $N_H$ and higher inclination angles. We note that the source with the lowest ratio of $L_{\\rm Fe}$/\\loiv~ is the prototypical Seyfert 2 galaxy NGC 1068, for which an edge-on orientation of the torus and a high, Compton-thick column density has already been suggested in literature (e.g., Pounds et al. 2006, and references therein)."
    },
    "1107/1107.5065_arXiv.txt": {
        "abstract": "We present new millimeter and radio observations of nine z$\\sim$6 quasars discovered in deep optical and near-infrared surveys. We observed the 250 GHz continuum in eight of the nine objects and detected three of them. New 1.4 GHz radio continuum data have been obtained for four sources, and one has been detected. We searched for molecular CO (6-5) line emission in the three 250 GHz detections and detected two of them. Combining with previous millimeter and radio observations, we study the FIR and radio emission and quasar-host galaxy evolution with a sample of 18 z$\\sim$6 quasars that are faint at UV and optical wavelengths (rest-frame 1450 $\\rm \\AA$ magnitudes of $\\rm m_{1450}\\geq20.2$). The average FIR-to-AGN UV luminosity ratio of this faint quasar sample is about two times higher than that of the bright quasars at z$\\sim$6 ($\\rm m_{1450}<20.2$). A fit to the average FIR and AGN bolometric luminosities of both the UV/optically faint and bright z$\\sim$6 quasars, and the average luminosities of samples of submillimeter/millimeter-observed quasars at z$\\sim$2 to 5, yields a relationship of $\\rm L_{FIR}\\sim {L_{bol}}^{0.62}$. Five of the 18 faint z$\\sim$6 quasars have been detected at 250 GHz. These 250 GHz detections, as well as most of the millimeter-detected optically bright z$\\sim$6 quasars, follow a shallower trend of $\\rm L_{FIR}\\sim {L_{bol}}^{0.45}$ defined by the starburst-AGN systems in local and high-z universe. The millimeter continuum detections in the five objects and molecular CO detections in three of them reveal a few $\\rm \\times10^{8}\\,M_{\\odot}$ of FIR-emitting warm dust and $\\rm 10^{10}\\,M_{\\odot}$ of molecular gas in the quasar host galaxies. All these results argue for massive star formation in the quasar host galaxies, with estimated star formation rates of a few hundred $\\rm M_{\\odot}\\,yr^{-1}$. Additionally, the higher FIR-to-AGN luminosity ratio found in these 250 GHz-detected faint quasars also suggests a higher ratio between star formation rate and supermassive black hole accretion rate than the UV/optically most luminous quasars at z$\\sim$6. ",
        "introduction": "Understanding the formation and evolution of the first galaxies and the supermassive black holes (SMBHs) they host is a key goal of both observational and theoretical astronomy. The first sample of more than twenty bright quasars at z$\\sim$6 were selected from the imaging data of the Sloan Digital Sky Survey (hereafter SDSS main survey, e.g., Fan et al. 2000, 2006). These objects are very bright in quasar UV and optical emission (SDSS magnitudes of $\\rm 18.74\\leq z_{AB}\\leq20.42$) and have SMBHs with masses of $\\rm M_{BH}\\sim$ a few $\\rm10^{9}\\,M_{\\odot}$ (Jiang et al. 2007; Kurk et al. 2007). The dust and molecular gas properties of this bright quasar sample at z$\\sim$6 have been studied at millimeter wavelengths, using the Max Planck Millimeter Bolometer Array (MAMBO), the IRAM Plateau de Bure Interferometer (PdBI), and the Very Large Array (VLA, Bertoldi et al. 2003a; 2003b; Petric et al. 2003; Walter et al. 2003, 2004, 2009; Carilli et al. 2004, 2007; Riechers et al. 2009; Wang et al. 2007, 2008b, 2010). About 30\\% of them have strong continuum emission at 250 GHz from 40-to-60 K warm dust (Bertoldi et al. 2003a; Petric et al. 2003; Wang et al. 2007, 2008b). The FIR luminosities of the MAMBO detections are a few $\\rm \\times10^{12}\\,L_{\\odot}$ to $\\rm 10^{13}\\,L_{\\odot}$, and the estimated dust masses are a few $\\rm \\times10^{8}\\,M_{\\odot}$ (Beelen et al. 2006; Bertoldi et al. 2003a; Wang et al. 2008a). Most of the 250 GHz-detected sources have also been detected in molecular CO line emission with PdBI, yielding molecular gas masses of $\\rm (0.7-3)\\times10^{10}\\,M_{\\odot}$ (Walter et al. 2003; Carilli et al. 2007; Riechers et al. 2009; Wang et al. 2010, 2011). The FIR-CO and FIR-radio luminosity ratios of these objects are consistent with the trend defined by dusty starburst systems at low and high redshifts (Riechers et al. 2006; Beelen et al. 2006; Wang et al. 2008b). The results provide direct evidence for massive star formation at a rate of $\\rm \\sim10^{3}\\,M_{\\odot}\\,yr^{-1}$, indicating that bulges are being formed at the same time that the SMBH is rapidly accreting in at least 1/3 of the optically bright quasar systems at z$\\sim$6 (Wang et al. 2008b). Molecular CO (3-2) and (7-6) line emission was resolved in one of the FIR and CO luminous quasars, SDSS J114816.64+525150.3 at z=6.42, revealing a source size of $\\rm \\sim5\\,kpc$ (Walter et al. 2004, 2009). The dynamical mass estimated with the size and line width of the CO emission suggest a SMBH-spheroidal bulge mass ratio an order of magnitude higher than the typical present-day value (Walter et al. 2004).  Z$\\sim$6 quasars that are fainter at UV and optical wavelengths have been discovered from deep optical-near infrared surveys (the SDSS southern deep imaging survey, $\\rm 20.54\\leq z_{AB}\\leq22.16$, Jiang et al. 2008, 2009, the Canada-France High redshift Quasar Survey [CFHQS], $\\rm 20.26\\leq z_{AB}\\leq24.40$, Willott et al. 2007, 2009a, 2009b). Near-infrared spectroscopic observations of some of these faint quasars reveal that they are likely to be accreting at the Eddington rate, and have less massive black holes (i.e., 10$\\rm ^{8}$ to $\\rm 10^{9}\\,M_{\\odot}$) than the SDSS main survey z$\\sim$6 quasar sample (Kurk et al. 2009; Willott et al. 2010b). The steepness of the quasar luminosity function at high redshifts is such that lower-luminosity/mass objects are more common. Thus, these faint z$\\sim$6 quasars provide an important sample to understand SMBH-host galaxy formation at the earliest epoch.  The millimeter dust continuum emission from the faint z$\\sim$6 quasars was first studied with a small sample of four objects discovered by CFHQS (Willott et al. 2007). One of them was detected at 250 GHz with a flux density of $\\rm \\sim1\\,mJy$, indicating a FIR continuum luminosity from dust of $\\rm \\gtrsim10^{12}\\,L_{\\odot}$. The mean 250 GHz flux density of the four quasars when stacked is $\\rm \\sim 0.6\\,mJy$, which is about a factor of two smaller than the average emission of the z$\\sim$6 luminous quasar sample (Willott et al. 2007; Wang et al. 2008b). We subsequently detected 250 GHz dust continuum and molecular CO (6-5) line emission in another two z$\\sim$6 quasars with $\\rm z_{AB}\\sim20.7$ (SDSS J205406.42-000514.8 and NDWFS J142516.30+325409.0, Jiang et al. 2008; Cool et al. 2006; Wang et al. 2008b, 2010) and CO (2-1) line emission from one CFHQS quasar at z=6.2 (CFHQS J142952+544717\\footnote{We exclude this object from our analysis throughout this paper as there is no published millimeter continuum observation yet.}, Willott et al. 2010a; Wang et al. 2011). These observations revealed a few $\\rm 10^{8}\\,M_{\\odot}$ of 40-to-60 K warm dust and $\\rm 10^{10}\\,M_{\\odot}$ of molecular gas in the quasar host galaxies, i.e., masses comparable to those found with the bright z$\\sim$6 quasars.  In this paper, we report new millimeter and radio observations of nine z$\\sim$6 quasars. We then study the FIR and radio emission with a sample of 18 z$\\sim$6 quasars that are faint at UV and optical wavelengths, and investigate the dust and molecular gas emission properties and star forming activity in the host galaxies of the millimeter-detected objects. We summarize the current quasar sample at z$\\sim$6 and describe the new millimeter and radio observations in Section 2, present the results in Section 3, analyse the continuum emission properties of the faint z$\\sim$6 quasar sample in Section 4, discuss the star formation and quasar-host evolution in Section 5, and conclude in Section 6. A $\\rm \\Lambda$-CDM cosmology with $\\rm H_{0}=71km\\ s^{-1}\\ Mpc^{-1}$, $\\rm \\Omega_{M}=0.27$ and $\\rm \\Omega_{\\Lambda}=0.73$ is adopted throughout this paper (Spergel et al. 2007). ",
        "conclusions": "We present observations of millimeter and radio continuum and CO(6-5) line emission from the host galaxies of quasars at z$\\sim$6. The new observations complete our MAMBO dust continuum survey of all the UV/optically faint z$\\sim$6 quasars ($\\rm m_{1450}\\geq20.2$) discovered from SDSS. Combining with previous data, we calculate the average FIR and radio emission for a sample of 18 z$\\sim$6 quasars with $\\rm m_{1450}\\geq20.2$. The mean FIR-to-AGN UV luminosity ratio of this faint quasar sample is about two times higher than that of the bright z$\\sim$6 quasars ($\\rm m_{1450}<20.2$). A fit to the average FIR and AGN bolometric luminosities of both the UV/optically faint and bright z$\\sim$6 quasars, and the average luminosities of samples of submillimeter/millimeter-observed quasars at z$\\sim$2 to 5, yields a relationship of $\\rm L_{FIR}\\sim {L_{bol}}^{0.62}$. The millimeter observations have detected the most FIR luminous objects among these faint z$\\sim$6 quasars. The 250 GHz dust continuum detections in five of the 18 faint quasars and CO (6-5) line detections in three of them (see Table 1 and 2) reveal dust masses of a few $\\rm 10^{8}\\,M_{\\odot}$ and molecular gas masses of $\\rm 10^{10}\\,M_{\\odot}$ in the quasar host galaxies, which are comparable to the dust and gas masses found in the bright z$\\sim$6 quasars (Beelen et al. 2006; Bertoldi et al. 2003b; Carilli et al. 2007; Wang et al. 2008a, 2010). Their FIR luminosities are estimated to be $\\rm 2.6\\sim5.5\\times10^{12}\\,L_{\\odot}$, and the FIR and AGN bolometric luminosities follow a shallower relationship of $\\rm L_{FIR}\\sim {L_{bol}}^{0.45}$, in agreement with the starburst-AGN systems at lower redshifts. All these results suggest active star formation in the host galaxies of the millimeter-detected and UV/optically faint quasars at z$\\sim$6 with SFRs of a few hundred $\\rm M_{\\odot}\\,yr^{-1}$. Moreover, the higher average FIR-to AGN UV luminosity ratios (Table 4) found with these objects and the shallow luminosity relationship suggest higher SFR-to-AGN accretion rate ratios than that of the more luminous/massive z$\\sim$6 quasars. Further high-resolution imaging of the dust and molecular gas components (e.g., with ALMA) in the quasar host galaxies will be critical to constrain the host galaxy dynamical mass, black hole-spheroidal host mass ratio, gas and star formation rate surface density of these objects."
    },
    "1107/1107.6017.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {Recently it has been reported that the intrinsic dispersion at constant magnitude of the structural relations from early-type galaxies is a useful tool to study the universality of these structural relations, that is to say, to study whether the structural relations depend on luminosity, wavelength, redshift and/or environment. In this work we study the intrinsic dispersion of the Faber-Jackson relation as function of the luminosity, mass and redshift.} % methods heading (mandatory) {We use a sample of approximately 90 000 early-type galaxies from the Sloan Digital Sky Survey (SDSS-DR7) spanning a magnitude range of 7 $mag$ in both $g$ and $r$ filters. We calculate the intrinsic dispersion of the Faber-Jackson relation at approximately constant magnitude and compare this at different luminosities, masses and redshifts.} % results heading (mandatory) {The main results are the following:  i) The intrinsic dispersion of the Faber-Jackson relation depends on the luminosity, mass and redshift. ii) The distribution for brighter and more massive galaxies has smaller intrinsic dispersion than that for fainter and less massive galaxies. iii) The distribution of bright and massive galaxies at higher redshift has smaller intrinsic dispersion than those similar galaxies at low redshift. } % conclusions heading (optional), leave it empty if necessary {Comparisons of the results found in this work with recent studies from the literature make us conclude that the intrinsic dispersion of the Faber-Jackson relation could depend on the history of galaxies, in other words, the intrinsic dispersion could depend on the number and nature of transformation events that have affected the galaxies along their life times, such as collapse, accretion, interaction and merging. } ",
        "introduction": "The scale relations of early-type galaxies (ETGs) are mathematical relations involving the galaxies' structural parameters such as: effective radius ($r_{e}$), mean effective surface brightness inside $r_{e}$ ($\\left\\langle \\mu\\right\\rangle _{e}$), central velocity dispersion ($\\sigma_0$) and total absolute magnitude ($M$). Among these structural relations, we have the Kormendy relation (KR; \\cite{kor77}), the Faber-Jackson relation (FJR; \\cite{fab76}) and the Fundamental Plane (FP; \\cite{djo87}; \\cite{dre87}), which are very useful tools for understanding the processes of formation and evolution of galaxies. In the recent past there have been many efforts directed to investigate the properties of ETGs using the structural relations. The large majority of these works have studied the coefficients of the structural relations at different wavelengths, environments, redshifts and luminosities and the results that they have found are very heterogeneous and discrepant. For example, several studies have found that the coefficients of the structural relations remain stable when considering different wavelengths (\\cite{ben92}; \\cite{BENDERETAL98}; \\cite{ber03b}; \\cite{ber03c}; \\cite{lab05}; \\cite{lab08}). However, other studies show that the structural relations depend on wavelength (\\cite{jor96}; \\cite{hud97}; \\cite{pah98}; \\cite{sco98}; \\cite{jun08}). Several studies have shown that the structural relations and/or the structural parameters of galaxies are affected by the environment (\\cite{ben92}; \\cite{BENDERETAL98}; \\cite{tru01}; \\cite{tru02}; \\cite{ber03b}; \\cite{agu04}; \\cite{gut04}; \\cite{den05}; \\cite{jor05}), although other studies have found the opposite conclusion (\\cite{ros01}; \\cite{tre01}; \\cite{evs02}; \\cite{gon03}; \\cite{red04}; \\cite{red05}; \\cite{nig07}). When structural relations of galaxy samples are studied at different redshifts, several studies indicate that only the zero point of these relations depends on the redshift (\\cite{bar98}; \\cite{zie99}; \\cite{lab03}; \\cite{bar06}). However, other authors find that there is, not only a dependence of the zero point on redshift but that the slopes of the structural relations are steeper for higher redshift galaxies than for galaxies in the local Universe (\\cite{tre05}; \\cite{jor06}; \\cite{fri09}). Finally, and regarding the luminosity, some authors find that dwarf and bright ellipticals follow structural relations with different coefficients (\\cite{kor85}; \\cite{ham87}; \\cite{ben92}; \\cite{cao93}; \\cite{agu09}; \\cite{des07}). Studies by Nigoche-Netro (2007) and Nigoche-Netro et al. (2008;2009;2010) find that the distribution of galaxies in the space of parameters that define the structural relations depends on the luminosity and that the coefficients of the structural relations depend on the width of the magnitude range ($\\Delta M$) within which the galaxies are contained. They also find that when the width of the magnitude range diminishes, the differences in the slope of the structural relations (for intervals of the same width and different luminosity) become small and when it has approximately constant magnitude the differences are negligible. This effect is present in all samples of galaxies studied, independently of their degree of intrinsic difference, and has been denoted as geometrical effect.   On the other hand, the analysis of the intrinsic dispersion is less frequent in the astronomical literature than that of the values of the coefficients of the structural relations. Some papers have found that this dispersion depends both on luminosity (\\cite{ben92}; \\cite{jor96}; \\cite{hyd09}) and the environment (\\cite{ber03b}; \\cite{den05}). The behaviour of the intrinsic dispersion of the FJR and the KR have not been studied thoroughly because it is considered that there is a third variable which causes most of the intrinsic dispersion in both these relations. The intrinsic dispersion of the FP has not been thoroughly studied because, it has been traditionally thought small ($\\sim 0.1$ dex, \\cite{kja93}, \\cite{jor96}, \\cite{kel97},  \\cite{jor99}, \\cite{bla02}, \\cite{ber03c}, \\cite{red05}, \\cite{jor06}).  However, some studies show that the intrinsic dispersion values are far from being small ($\\sim0.3$ dex, \\cite{ben92}; \\cite{lab03}; \\cite{nig09}). This, relatively high dispersion causes that the galaxy distribution in the space that defines the structural relations follows a surface whose thickness is determined by this dispersion.  Recent works (\\cite{nig07b}; Nigoche-Netro et al. 2008; 2009; 2010) have shown that the intrinsic dispersion is also affected by the geometrical effect. However, Nigoche-Netro et al. (2010) have found that differences in the value of the intrinsic dispersion for different samples of galaxies do not disappear as the magnitude range diminishes.  They find that the exact value for the intrinsic dispersion is obtained when $\\Delta M = 0$. The intrinsic dispersion for this extreme case would be defined as the standard deviation of the distribution of the points at constant magnitude. Hence, Nigoche-Netro et al. (2010) consider that an appropriate method for obtaining physical information for a sample of galaxies is to find its intrinsic dispersion at each magnitude value and then perform comparisons of this dispersion at different luminosities, wavelengths, redshifts, or environments. In light of this, in this paper we carry out a study of the behaviour of the intrinsic dispersion of the FJR as a function of luminosity, mass and redshift for a sample of ETGs selected from the SDSS-DR7 archive.  In \\S~2 we present the galaxy sample used to study the intrinsic dispersion of the FJR, the calculation of the intrinsic dispersion and the analysis of the behaviour of the intrinsic dispersion as a function of the luminosity, mass and redshift. In \\S~3 we present a discussion of the most important results of this paper. Finally in \\S~4 we present our conclusions. ",
        "conclusions": "Analysing the intrinsic dispersion of the FJR for a sample of approximately 90 000 ETGs from the literature we find the following:  \\begin{itemize}  \\item The intrinsic dispersion of FJR depends on the luminosity, mass and redshift.  \\item The distribution of brighter and more massive galaxies has a lower intrinsic dispersion than the distribution for the fainter and less massive galaxies.  \\item In the redshift range $0.01 < z < 0.25$, for luminous and massive galaxies, the distribution of the farther galaxies has a lower intrinsic dispersion than that for the nearer galaxies.   \\end{itemize}   A recent work (\\cite{fra10}) has shown that the structural relations, among which we have the FJR, are formed by 7 different groups of galaxies. These groups define separate regions on the structural relations. Each group defines its own structural relation which is more loosely defined for less-diversified groups. The `diversity' means in this context the number and the nature of transformation events that affect the galaxies along their life times. When we compare our results with those of Fraix-Burnet et al. (2010) we find that the distribution of galaxies inside the groups and the distribution of the groups on the FJR plane might cause that the brighter (more-diversified) galaxies have lower intrinsic dispersion than the fainter (less-diversified) galaxies. So we may conclude that the intrinsic dispersion of the FJR might depend on the number and the nature of transformation events that affect the galaxies such as collapse, accretion, interaction and merging, that is to say, the intrinsic dispersion might depend on the histories of galaxies.  As mentioned above (Section 3), before taking the former results as conclusive, further diversification studies should be carried out in ETGs samples similar to the one presented in this paper. It is also important to consider that the intrinsic dispersion might depend on other factors which have not been considered in the analysis by Fraix-Burnet et al. (2010) such as the fact that part of this dispersion might be due to original inherent properties of the galaxies and not only to the transformation events which they have suffered along their life times.  Finally, as seen in section 2.3, a comparison of virial and stellar mass returns a value of the slope approximately equal to 1, this value differs from recent results in the literature in which the slope could be as large as $\\sim 1.3$. This would mean that inside the virial radius 30\\% of the mass could be in the form of dark matter. Our results indicate that inside this radius the amount of dark matter is negligible. We shall discuss this point in a forthcoming paper."
    },
    "1107/1107.3643_arXiv.txt": {
        "abstract": "A complete census of planetary systems around a volume-limited sample of solar-type stars (FGK dwarfs) in the Solar neighborhood ($d\\leq15\\,pc)$ with uniform sensitivity down to Earth-mass planets within their Habitable Zones out to several AUs would be a major milestone in extrasolar planets astrophysics.  This fundamental goal can be achieved with a mission concept such as NEAT --- the Nearby Earth Astrometric Telescope.  NEAT is designed to carry out space-borne extremely-high-preci\\-sion astrometric measurements at the $0.05\\,\\mu$as ($1 \\sigma$) accuracy level, sufficient to detect dynamical effects due to orbiting planets of mass even lower than Earth's around the nearest stars. Such a survey mission would provide the actual planetary masses and the full orbital geometry for all the components of the detected planetary systems down to the Earth-mass limit. The NEAT performance limits can be achieved by carrying out differential astrometry between the targets and a set of suitable reference stars in the field. The NEAT instrument design consists of an off-axis parabola single-mirror telescope (D = 1m), a detector with a large field of view located 40\\,m away from the telescope and made of 8 small movable CCDs located around a fixed central CCD, and an interferometric calibration system monitoring dynamical Young's fringes originating from metrology fibers located at the primary mirror. The mission profile is driven by the fact that the two main modules of the payload, the telescope and the focal plane, must be located 40\\,m away leading to the choice of a formation flying option as the reference mission, and of a deployable boom option as an alternative choice. The proposed mission architecture relies on the use of two satellites, of about 700 kg each, operating at L2 for 5 years, flying in formation and offering a capability of more than 20,000 reconfigurations. The two satellites will be launched in a stacked configuration using a Soyuz ST launch vehicle.  The NEAT primary science program will encompass an astrometric survey of our 200 closest F-, G- and K-type stellar neighbors, with an average of 50 visits each distributed over the nominal mission duration. The main survey operation will use approximately 70\\% of the mission lifetime.  The remaining ~30\\% of NEAT observing time might be allocated, for example, to improve the characterization of the architecture of selected planetary systems around nearby targets of specific interest (low-mass stars, young stars, etc.) discovered by Gaia, ground-based high-precision radial-velocity surveys, and other programs. With its exquisite, surgical astrometric precision, NEAT holds the promise to provide the first thorough census for Earth-mass planets around stars in the immediate vicinity of our Sun. ",
        "introduction": "\\label{sec:introduction}  Exoplanet research has grown explosively in the past decade, supported by improvements in observational techniques that have led to increasingly sensitive detection and characterization. Among many results, we have learned that planets are common, but their physical and orbital properties are much more diverse than originally thought.  A lasting challenge is the detection and characterization of planetary systems consisting in a mixed cortege of telluric and giant planets, with a special regard to telluric planets orbiting in the habitable zone (HZ) of Sun-like stars. The accomplishment of this goal requires the development of a new generation of facilities, due to the intrinsic difficulty of detecting Earth-like planets with existing instruments. The proposed NEAT mission has been designed to enter a new phase in exoplanetary science by delivering an enhanced capability of detecting small planets at and beyond 1\\,AU.  Astrometry is probably the oldest branch of astronomy. Greeks developed it and noticed that the position of most stars were stable in the sky, but the few that were moving became known as \\emph{planets} (\\textgreek{pl'anetes >ast'eres} = moving stars), pointing to a major difference in their nature. Thanks to the precise astrometric measurements of planet positions by Tycho Brahe in the 16th century, Johannes Kepler established that these objects were orbiting the Sun on elliptical orbits, expanding the Copernican revolution. After Hipparcos, Gaia will play an important role in finding many systems with giant planets in our Galaxy. We want to extend these revolutions with the NEAT mission, namely to discover and characterize Earth-mass planets in Earth-like orbits around stars like the Sun, by capturing infinitesimal displacements with unprecedented accuracy.  In Sect.~\\ref{sec:neat-science}, we present the science objectives of NEAT, we describe the principle of the differential astrometry technique and we give a list of potential targets. In Sect.~\\ref{sec:neat-concept}, after listing the technical challenges, we present the instrumental concept. We explain how to reach the performance and we give a summarized description of the payload, the mission and the spacecraft. In Sect.~\\ref{sec:discussion}, we discuss both astrophysical and technical issues. Recommandations by the community summarized in Sect.~\\ref{sec:perspectives} is an incentive to pursue the development of this mission in the future. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Astrophysical issues} \\label{sec:astro-challenges}  \\textbf{Stellar activity.} If all instrumental problems are controlled then the next obstacle to achieve the scientific objective is of astrophysical nature, the impact of stellar activity. Spots and bright structures on the stellar surface induce astrometric, photometric and RV signals. Using the Sun as a proxy, \\citet{2011A&A...528L...9L} have computed the astrometric, photometric and RV variations that would be measured from an observer located 10 pc away. It appears that the astrometric variations due to spots and bright structures are small compared to the signal of an Earth mass planet in the HZ \\citep[see also][]{2010A&A...519A..66M, 2009ApJ...707L..73M, 2010ApJ...717.1202M}. This remains true throughout the entire solar cycle. If we consider a star \\emph{5 times more active than the active Sun}, an Earth-mass planet would still be detectable even during the highest activity phases. Such activity, or lower, translates in terms of activity index $\\log(R'_{HK}) \\leq -4.35$. Consequently, in our target list, we have kept only stars with such an index (\\emph{only 4\\% were discarded}), for which their intrinsic activity should not prevent the detection of an Earth-mass planet, even during its high activity period.  \\textbf{Perturbations from reference stars.} The vast majority of the reference stars will be K giants at a distance of $\\approx 1$\\,kpc. The important parameters in addition to the position are the proper motion with typical value of $\\approx 1$\\,mas/yr and the parallax whose typical value is $\\approx 1$\\,mas.  They are to be compared to the accuracy of the cumulative measurements during a visit. An important value for NEAT accuracy is what is obtained for an $R=6$ magnitude target: 0.8\\,\\uas/h. The ratios between that (required) accuracy and the expected motions of the references indicates clearly that the latter cannot be considered as fixed. Their positions are members of the set of parameters that have to be solved for. Because the reference stars are much more distant ($\\approx1$\\,kpc) than the target star ($\\approx10$\\,pc), we are 100 times less sensitive to their planetary perturbations. Only Saturn-Jupiter mass objects matter, and statistically, they are only present around $\\approx10$\\% of stars. These massive planets can be searched for by fitting first the reference star system ($\\approx 100 N_{\\rm ref}$ measurements for $5 N_{\\rm ref}$ parameters when there are no giant planets around the reference stars), possibly eliminate those with giant planets, and studying the target star with respect to that new reference frame. Moreover, the largest disturbers will be detected from ground based radial velocity measurements, and the early release of Gaia data around 2016 will greatly improve the position accuracy of the reference stars. For smaller planets at or below the threshold of detection, their impact on the target astrometry will be only at a level $\\ll 1\\,\\MEarth$ around it. Similarly the activity of these K giants has been investigated and neither the stellar pulsations nor the stellar spots will disturb the signal at the expected accuracy.  \\textbf{Planetary system extraction from astrometric data.} We recently carried out a major numerical simulation to test how well a space astrometry mission could detect planets in multi-planet systems \\citep{2010EAS....42..191T}.  The simulation engaged 5 teams of theorists who generated model systems, and 5 teams of double-blind \\emph{``observers''} who analyzed the simulated data with noise included.  The parameters of the study were the same as for NEAT, viz., astrometric single-measurement uncertainty (0.80\\,\\uas noise, 0.05\\,\\uas floor, 5-year mission, plus RV observations with 1\\,m/s accuracy for 15\\,years).  We found that terrestrial-mass, habitable-zone planets ($\\approx$ Earths) were detected with about the same efficiency whether they were alone in the system or if there were several other giant-mass, long-period planets ($\\approx$ Jupiters) present. The reason for this result is that signals with unique frequencies are well separated from each other, with little cross-talk.  The number of planets per system ranged from 1 to 11, with a median of 3. The SNR value of 5.8 value was predicted by \\citet{1982ApJ...263..835S} for a false alarm probability (FAP) of less than 1\\%, and verified in our simulations. The completeness and reliability to detect planets was better than 90\\% for all planets, where the comparison is with those planets that should have been detected according to a Cramer-Rao estimate \\citep{2010ApJ...720.1073G} of the mission noise.  The Cramer-Rao estimates of uncertainty in the parameters of mass, semi-major axis, inclination, and eccentricity were consistent with the \u201cobserved\u201d estimates of each: 3\\% for planet mass, $\\approx4^{\\circ}$ for inclination and 0.02 for eccentricity.  \\textbf{Radial velocity screening.} To solve unambiguously for giant planets with periods longer than 5\\,yrs, it is necessary to have a ground RV survey for 15\\,yrs of the 200 selected target star, at the presently available accuracy of 1\\,m/s. More than 80\\% of our targets are already being observed by RV, but the observations of the rest of them should start soon, well before the whole NEAT data is available.  The capability of ground based RV surveys, despite their impressive near-term potential to obtain accuracies better than 1\\,m/s, is not sufficient to detect terrestrial planets in the HZ of F, G and K stars.  Formally, an accuracy of 0.05\\,m/s is required to see an edge-on Earth mass planet at 1\\,\\AU from a solar-mass star with SNR=5 (semi-amplitude = 0.13\\,m/s), which might be achievable instrumentally, but is stopped in most cases by the impact of stellar activity on RV accuracy. It is necessary to find particularly ``quiet'' stars, but they are a minority (few percents) and cannot provide a full sample. Furthermore, the ambiguity in physical mass associated with the signal coming only from the radial component of the stellar reflex motion ($\\sin i$ ambiguity) requires additional information to determine the physical mass and relative inclination in complex planetary systems. In some, but not all cases, limits are possible, and one can argue statistically that 90\\% of systems should be oriented such that the physical planet mass is within a factor of two of the mass found in RV. However, for finding a small number of potential future targets for direct detection and spectroscopy, an absolute determination that the mass is Earth-like is required as well as an exhaustive inventory of the planets around stars in our neighborhood.  \\begin{table*}[t] \\centering \\caption{Science impact of NEAT scaling. The nominal mission is highlighted in yellow.} \\label{tab:scaling} \\includegraphics[trim=0mm 7mm 5mm 5mm,clip, width=0.8\\textwidth]{neat-scaling-table} \\end{table*} \\textbf{Flexibility of objectives to upgrades / downgrades of the mission.} One of the strengths of NEAT is its flexibility, the possibility to adjust the size of the instrument with impacts on the science that are not prohibitive. The size of the NEAT mission could be reduced (or increased) with a direct impact on the accessible number of targets but not in an abrupt way. For instance, for same amount of integration time and number of maneuvers, the options listed in Table \\ref{tab:scaling} are possible, with impacts on the number of stars that can be investigated down to 0.5 and 1 Earth mass, and on the mass of the instrument, required fuel for maneuvers, and therefore cost. The time necessary to achieve a given precision depends on the mass limit that we want to reach: going from 0.5\\,\\MEarth to 1\\,\\MEarth requires twice less precision and therefore 4 times less observing time allowing a smaller telescope. There is room for adjustment keeping in mind that one wants to survey the neighborhood with the smallest mass limit possible and a typical number of targets of $\\approx 200$.   \\subsection{Technical issues} \\label{sec:technical-challenges}  \\textbf{Optical aberrations.} NEAT uses a very simple telescope optical design. A 1-m diameter clear aperture off-axis parabola, with an off-axis distance of 1\\,m and a 40\\,m focal length. The focal plane is at the prime focus. The telescope is diffraction limited at the center of the field, where the target stars will be observed, but coma produces some field dependent aberrations. At the mean position of the reference stars, $0.2^{\\circ}$ away from the center of the field, the coma produces a \\emph{steady} 23\\% increase of the point spread function (PSF) width and an 8\\,\\microns centroid offset. The impact remains low since we are looking at differential effects.  \\textbf{Centroid measurements.} They consist of two steps: the determination of the stellar centroid on each CCD during 57\\,s and then the calibration of the relative position of the CCDs during 3\\,s thanks to the metrology. The metrology determines also the response map of the detectors. As in the normal approach to precision astrometry with CCDs, we perform a least-square fit of a template PSF to the pixelated data. PSF knowledge error leads to systematic errors in the conventional centroid estimation. We have developed an accurate centroid estimation algorithm by reconstructing the PSF from well sampled (above Nyquist frequency) pixelated images. In the limit of an ideal focal plane array whose pixels have identical response function (no inter-pixel variation), this method can estimate centroid displacement between two 32x32 images to sub-micropixel accuracy. Inter-pixel response variations exist in real CCDs, which we calibrate by measuring the pixel response of each pixel in Fourier space\\footnote{They are determined by calculating the first 6 coefficients of the Taylor series expansion in powers of wave numbers of the detector response map Fourier components.}. Capturing inter-pixel variations of pixel response to the third order terms in the power series expansion, we have shown with simulated data that the centroid displacement estimation is accurate to a few micro-pixels.  \\textbf{Stability of the primary mirror.} The primary optic will be made of zerodur/ULE with a temperature coefficient better than $10^{-8}$/K with an optics thickness $\\approx 10$\\,cm and the effective temperature and temperature gradients are kept stable to $\\approx 0.1$\\,K over the mirror, the optic is then stable to $\\approx0.1$\\,nm ($\\lambda/6000$) during the 5\\,yr mission. We have simulated two images, one at the center of the field that is a perfect Airy function and one at the edge of the field that has a $\\lambda/20$ coma. We added also wavefront errors with a conservative rms value of $\\lambda/1000$. With the new wavefronts, we calculated the change in the differential astrometry bias caused by both pixelation and changing wavefronts. While the wavefront deviations to optimal shape caused a centroid shift of $\\approx6-10\\,\\uas$ ($10^{-4}$ pixels), differential errors remained less than $\\approx0.3\\,\\uas$ ($3\\times10^{-6}$ pixels).  \\textbf{CCD damage in L2 environment.} CCDs, like most semiconductors, suffer damage in radiation environments such as encountered by space missions. One particular performance parameter, Charge Transfer Efficiency (CTE), degrades with known consequences on the efficiency of science missions like Gaia\\footnote{The Gaia community (\\url{http://www.rssd.esa.int/gaia}) speaks of the complementary quantity, charge transfer inefficiency (CTI), in order to emphasize its detrimental effects.}. The reduced CTE is caused generally by prompt particle events (PPE), including solar protons and cosmic rays, colliding with the CCD silicon lattice and causing damages to the silicon lattice. This leads to the formation of so-called traps which can capture photo-electrons and release them again after some time. This results in signal loss and distortion of the PSF shape. The latter leads to systematic errors in the image location due to a mismatch between the ideal PSF shape and the actual image shape. For Gaia, this effect of radiation damage is a major contributor to the error budget and extensive research and laboratory tests have been done in order to understand better the radiation damage effects and to develop approaches in both hardware and data processing to mitigate the negative impact. However, there are a number of important differences between NEAT and Gaia which justify the assumption that radiation damage effects will play a much smaller role: i) NEAT looks for extended periods at very bright stars compared to Gaia in which the stars continuously move on the CCD. Also, unlike Gaia, NEAT will not be operated in time-delayed integration mode. In addition the CCDs are regularly illuminated by the laser light from the metrology system. This means that in general the signal level in the CCD pixels is high which will keep the traps with long ($\\geq60$\\,s) release time constants filled and effectively inactive. ii) NEAT also does not suffer from the varying CCD illumination history that a scanning mission like Gaia necessarily encounters. This illumination history is in fact one of the major complicating factors for Gaia. Finally, iii) NEAT uses much smaller CCDs than Gaia and in addition has four read-out nodes, thus reducing the number of charge transfer steps and mitigating the effects of radiation damage. The one concern for the NEAT case is the presence of traps with release time constants that are of the order of several times the charge transfer period between pixels. In the case of NEAT the transfer period averages tens of $\\mu$s and from laboratory tests with E2V CCDs, carried out in the context of the Gaia project, traps with time constants of 10--100\\,$\\mu$s are known to exist. If these traps dominate at the operating temperature of the NEAT CCDs they could lead to subtle PSF image shape distortions and thus image location biases. From the Gaia experience, it is known that such image shape distortions can be handled in the post-processing by a careful modeling of the effects of radiation damage on the PSF image. A similar strategy, building on the Gaia heritage, can be employed for NEAT.   \\begin{figure*}[t] \\centering \\includegraphics[width=\\textwidth]{testbed-results} \\caption{Latest results obtained from the CCD/metrology test bench. No metrology nor quantum efficiency 6-parameter calibration have been performed yet. Left: Allen deviation of the centroid location for one artificial star ``A'' projected on the CCD and Allen deviation of the differential centroid location for two artificial stars A and B projected on the CCD. One can see that star A moves on the CCD by a few hundred micro-pixels at time scale greater than 1 second, but that the differential position of the two stars is better $4\\,10^{-5}$ pixel at 100\\,s integration time. Right: Allen deviation of the differential centroid location for two artificial stars projected on the CCD. Concatenating data from 38 runs, a minimum of $\\approx 20\\,\\mu$pixels at about 10\\,min before differential drift dominated.  This data shows that differential metrology at intervals of minutes is needed.} \\label{fig:testbed-results} \\end{figure*} \\textbf{CCD/metrology tests in the lab.} In the absence of optical errors, the major error sources are associated with the focal plane: (1) motions of the CCD pixels, which have to be monitored to $3\\times10^{-6}$ pixels every 60\\,s, i.e.\\ 0.03\\,nm; (2) measurements of the centroid of the star images with $5\\times10^{-6}$ pixel accuracy. We have set up technology testbeds to demonstrate that we can achieve these objectives. The technology objective for (1) has almost been reached and the technology demonstration for (2) is underway and should be completed soon. Latest results with no metrology nor QE 6-parameter calibration have been obtained from the CCD / metrology test bench (Fig.~\\ref{fig:testbed-results}). Allen deviation of the centroid location for one artificial star ``A'' projected on the CCD and Allen deviation of the differential centroid location for two artificial stars A and B projected on the CCD are plotted in Fig.~\\ref{fig:testbed-results} (left). One can see that star A moves on the CCD by a few hundred micro-pixels at time scale greater than 1 second, but that the differential position of the two stars is better $4\\times10^{-5}$ pixel at 100s integration time. On Fig.~\\ref{fig:testbed-results} (right), the Allen deviation of the differential centroid location for two artificial stars projected on the CCD is plotted, concatenating data from 38 runs, a minimum of $\\approx20\\,\\mu$-pixels at about 10\\,min before differential drift dominated. This data shows that we are only a factor 10 from the final goal and that differential metrology at intervals of minutes is required to reach it."
    },
    "1107/1107.4063_arXiv.txt": {
        "abstract": "\\vspace*{2.5mm}\\noindent A particular compensation-type solution of the main cosmological constant problem has been proposed recently, with two massless vector fields dynamically canceling an arbitrary cosmological constant $\\Lambda$. The naive expectation is that such a compensation mechanism does not allow for the existence of an inflationary phase in the very early Universe. However, it is shown that certain initial boundary conditions on the vector fields can in fact give rise to an inflationary phase. ",
        "introduction": "\\label{sec:Intro}  Inflation~\\cite{Guth1981,Linde1983}, an epoch of exponential expansion, may have played an important role in the evolution of the very early universe (see Ref.~\\cite{precursor-papers} for an incomplete list of precursor papers and Ref.~\\cite{Mukhanov2005} for further references and discussion). The mechanism relies, however, on one crucial assumption [stated, for example, a few lines below Eq.~(3.7) in Ref.~\\cite{Guth1981}]: the minimum of the total scalar potential is set to zero, i.e., the corresponding vacuum energy density (effective cosmological constant $\\Lambda$) is assumed to vanish. In other words, it is taken for granted that a solution has been found to the main cosmological constant problem~\\cite{Weinberg1988} (CCP1): why is the present value of $|\\Lambda|^{1/4}$ negligible when compared with the known energy scales of elementary particle physics? The next cosmological constant problem (CCP2) is, of course, to explain the measured value $\\Lambda^\\text{exp} \\sim(2\\,\\text{meV})^4$, but this question lies outside the scope of the present article.  Following an earlier suggestion to consider vector fields~\\cite{Dolgov1985-1997} and using the insights from the $q$--theory approach~\\cite{KV2008-2010} to CCP1, a special model of two massless vector fields has been presented in Ref.~\\cite{EmelyanovKlinkhamer2011}. The massless vector fields of the model cancel dynamically an arbitrary initial (bare) cosmological constant $\\Lambda$ without upsetting the local Newtonian dynamics (a potential problem discussed in Ref.~\\cite{RubakovTinyakov1999}).  But, if any cosmological constant can be canceled dynamically, what happens to inflation in the very early universe?  In order to address this issue, we investigate the simplest possible extension of the two-vector-field model by adding a fundamental scalar field with a quadratic potential, while keeping an initial cosmological constant $\\Lambda$. The question is, then, whether or not it is possible to have an inflationary phase of finite duration. ",
        "conclusions": "\\label{sec:Discussion}  The results of Fig.~\\ref{fig:4new} were to be expected. The surprising (and encouraging?) results are those of Fig.~\\ref{fig:3new}, with an inflationary phase of some 10 $e$-foldings of $a(\\tau)$ for the model and parameters chosen. (Different initial conditions have been seen to give some 30 $e$-foldings and Fig.~\\ref{fig:2new} can be interpreted as having infinitely many $e$-foldings.) Qualitatively, the Hubble parameter $h(\\tau)$ of the top right panel in Fig.~\\ref{fig:3new} ($\\lambda \\ne 0$) resembles that of Fig.~\\ref{fig:1new} ($\\lambda = 0$), even though the detailed behavior differs as shown by the respective  bottom right panels.  The results of Fig.~\\ref{fig:3new} are, of course, only exploratory. It remains, for example, to analyze the nonlinear dynamics displayed in Fig.~\\ref{fig:3new} and to rigorously establish the $\\tau\\to\\infty$ limit for the $\\eta=10^{-4}$ case, corresponding respectively to an early phase with inflation\\footnote{Possible observable effects may come from the dynamics of density perturbations with near-horizon wavelengths (cf. Ref.~\\cite{Mukhanov2005}), as the dynamics can be expected to be modified by the interaction between the scalar (inflaton) field and the vector fields needed for the cancellation of the \\mbox{bare cosmological constant $\\Lambda$.}} and a late (FRW-like) phase with standard local Newtonian dynamics.  The main result of this article is that, in principle, it appears to be possible to have both an early phase with inflation and a late phase with a dynamically canceled cosmological constant $\\Lambda$. The details of the model are of secondary importance. What matters is the general mechanism which relies on the dynamics of the massless vector fields (or possibly massless tensor fields). The physical origin of these massless vector (tensor) fields needs to be clarified.  \\newpage \\begin{figure*}[th]% \\begin{center} \\hspace*{-6mm} \\includegraphics[width=1.05\\textwidth]{ccp1-inflation_fig1_v4.eps} \\end{center} \\vspace*{-5mm} \\caption{Numerical solution of the reduced field equations for model \\eqref{eq:model} and \\textit{Ansatz} \\eqref{eq:Dolgov-Ansatz}. The dimensionless model parameters are $m=0.01$, $\\delta=10^{-6}$, $\\eta=0$, and $\\lambda=0$. The boundary conditions are $a(1)=1$, $\\{\\varphi_1(1),\\,\\dot{\\varphi}_1(1),\\,\\varphi_2(1),\\,\\dot{\\varphi}_2(1)\\}$ $=$ $\\{10,\\, -0.25,\\, 0,\\, -0.433013 \\}$, and $v(1)=\\dot{v}(1)= w(1)=\\dot{w}(1)=0$. [The real scalars $\\varphi_n$ are defined by $\\sigma\\equiv (\\varphi_1+i\\,\\varphi_2)/\\sqrt{2}$.] The value of $h(1)$ follows from the Friedmann equation \\eqref{eq:Friedmann-ODE}. With these boundary conditions, the reduced field equations have the exact solution $v(\\tau)=w(\\tau)=0$ for $\\tau \\geq 1$. The top right panel shows an $h(\\tau)$ plateau corresponding to an inflationary phase. The bottom right panel shows that, long after inflation, $h(\\tau)$ asymptotically goes as $(2/3)\\,\\tau^{-1}$. \\vspace*{0cm} } \\label{fig:1new} \\end{figure*} \\begin{figure*}[h]% \\vspace*{0cm} \\begin{center} \\hspace*{-6mm} \\includegraphics[width=1.05\\textwidth]{ccp1-inflation_fig2_v4.eps} \\end{center} \\vspace*{-5mm} \\caption{Same as Fig.~\\ref{fig:1new}, but now with nonzero cosmological constant, $\\lambda=0.01$. Same boundary conditions as Fig.~\\ref{fig:1new}, so that the exact solution $v(\\tau)=w(\\tau)=0$ persists. The Friedmann equation \\eqref{eq:Friedmann-ODE} gives the asymptotic value $h(\\infty)=\\sqrt{\\lambda/3}\\approx 0.05774$. \\vspace*{-4cm} } \\label{fig:2new} \\end{figure*}  \\newpage \\begin{figure*}[th]% \\begin{center} \\hspace*{-6mm} \\includegraphics[width=1.05\\textwidth]{ccp1-inflation_fig3_v4.eps} \\end{center} \\vspace*{-5mm} \\caption{Same parameters and boundary conditions as Fig.~\\ref{fig:2new} (e.g., $\\lambda=0.01$), except for small but nonzero starting values of $v$ and $w$. Specifically, the boundary conditions at $\\tau=1$ are $\\{a,\\,v,\\,\\dot{v},\\,w,\\,\\dot{w},\\,h,\\, \\varphi_1,\\,\\dot{\\varphi}_1,\\,\\varphi_2,\\,\\dot{\\varphi}_2\\}$ $=$ $\\{1,\\,2\\times 10^{-4},\\,2\\times 10^{-4},\\,2\\times 10^{-4},\\,2\\times 10^{-4}$, $0.216025$, $10,\\, -0.25,\\, 0,\\, -0.433013 \\}$. The total vacuum energy density entering the right-hand side of the Friedmann equation \\eqref{eq:Friedmann-ODE} is given by the sum of the two panels in the middle column and vanishes asymptotically. \\vspace*{1\\baselineskip}} \\label{fig:3new} \\end{figure*} \\begin{figure*}[h]% \\vspace*{0cm} \\begin{center} \\hspace*{-6mm} \\includegraphics[width=1.05\\textwidth]{ccp1-inflation_fig4_v4.eps} \\end{center} \\vspace*{-5mm} \\caption{Same parameters and boundary conditions as Fig.~\\ref{fig:3new} (e.g., $\\lambda=0.01$), except for relatively large starting values of $v$ and $w$. Specifically, the boundary conditions at $\\tau=1$ are $\\{a,\\,v,\\,\\dot{v},\\,w,\\,\\dot{w},\\,h,\\, \\varphi_1,\\,\\dot{\\varphi}_1,\\,\\varphi_2,\\,\\dot{\\varphi}_2\\}$ $=$ $\\{1,\\,0.0799201,\\, 0.01998,\\, 0.08,\\, 0.02$, $0.216961$, $10,\\, -0.25,\\, 0,\\, -0.433013 \\}$. \\vspace*{-4cm} } \\label{fig:4new} \\end{figure*}   \\newpage%"
    },
    "1107/1107.4313_arXiv.txt": {
        "abstract": "The Small Magellanic Cloud (SMC) provides a unique laboratory for the study of the lifecycle of dust given its low metallicity ($\\sim$1/5~solar) and relative proximity ($\\sim$60~kpc).  This motivated the SAGE-SMC (Surveying the Agents of Galaxy Evolution in the Tidally-Stripped, Low Metallicity Small Magellanic Cloud) {\\em Spitzer} Legacy program with the specific goals of studying the amount and type of dust in the present interstellar medium, the sources of dust in the winds of evolved stars, and how much dust is consumed in star formation.  This program mapped the full SMC (30~$\\sq\\arcdeg$) including the Body, Wing, and Tail in 7 bands from 3.6 to 160~\\micron\\ using the IRAC and MIPS instruments on the {\\em Spitzer Space Telescope}. The data were reduced, mosaicked, and the point sources measured using customized routines specific for large surveys.  We have made the resulting mosaics and point source catalogs available to the community.  The infrared colors of the SMC are compared to those of other nearby galaxies and the 8~\\micron/24~\\micron\\ ratio is somewhat lower and the 70~\\micron/160~\\micron\\ ratio is somewhat higher than the average.  The global infrared spectral energy distribution shows that the SMC has approximately 1/3 the aromatic emission/PAH (polycyclic aromatic hydrocarbon) abundance of most nearby galaxies.  Infrared color-magnitude diagrams are given illustrating the distribution of different asymptotic giant branch stars and the locations of young stellar objects.  Finally, the average spectral energy distribution (SED) of HII/star formation regions is compared to the equivalent Large Magellanic Cloud average HII/star formation region SED.  These preliminary results will be expanded in detail in companion papers. ",
        "introduction": "\\label{sec_intro}  The interstellar medium (ISM) plays a central role in galaxy evolution as the birthsite of new stars and repository of old stellar ejecta. The formation of new stars slowly consumes the ISM, locking it up for millions to billions of years. As these stars age, the winds from low mass, asymptotic giant branch (AGB) stars, high mass, red supergiants (RSGs)/Wolf Rayet stars/Luminous Blue Variables (LBVs), and supernova explosions inject nucleosynthetic products of stellar interiors into the ISM, slowly increasing its metallicity. This constant recycling and associated enrichment drives the evolution of a galaxy's visible matter and changes its emission characteristics. To understand this recycling, we have to study the physical processes of the ISM, the formation of new stars, and the injection of mass by evolved stars, and their relationships on a galaxy-wide scale.  Among the nearby galaxies, the \\object{Small Magellanic Cloud} (SMC) represents a unique astrophysical laboratory for studies of the lifecycle of the ISM, because of its proximity \\citep[$\\sim$60 kpc,][]{Hilditch05}, low ISM metallicity \\citep[1/5--1/8 Z$_{\\sun}$,][]{Russell92, Rolleston99, Rolleston03, Lee05} and history of disruption by tidal interaction \\citep{Zaritsky00}.  The SMC offers a rare glimpse into the physical processes in an environment with a metallicity which is below the threshold of 1/4--1/3 Z$_{\\sun}$ where the properties of the ISM in galaxies changes significantly as traced by the rapid reduction in the aromatic emission/polycyclic aromatic hydrocarbon (PAH) dust mass fractions and dust-to-gas ratios \\citep{Engelbracht05SB, Draine07, Sandstrom10}.  In addition, the SMC is the only local galaxy which has the ultraviolet dust characteristics \\citep[lack of 2175~\\AA\\ extinction bump,][]{Gordon03} of starburst galaxies in the local \\citep{Calzetti94, Gordon97} and high-redshift \\citep[$2<z<4$,][]{Vijh03} universe.  The evolution of stars in the SMC is also clearly affected by their low metallicities \\citep{Cioni06, Marigo08} with the corresponding expected differences in stellar mass loss \\citep{vanLoon08}. The Large and Small Magellanic Clouds represent the nearest example of tidally interacting galaxies and the Magellanic Bridge is a clear manifestation of a close encounter of these two galaxies some 200 Myr ago \\citep{Zaritsky04, Harris07}.  Over cosmological timescales, galaxy interactions are one of the key drivers of galaxy evolution and, thus, tidally interacting galaxies allow us to examine star formation in an unusual and disturbed environment, which resembles the conditions in the early universe when galaxies were forming.  The Magellanic Bridge is a filament of neutral hydrogen, which joins the SMC and LMC over some 15 kpc \\citep{McGee86, Staveley-Smith98, Muller03}. Recent studies have revealed the presence of young ($<$200 Myrs) massive stars \\citep{Harris07} associated with the highest-density portion of the Bridge directly adjacent to the SMC main body.  Given the relative youth of these stars, they are highly likely to have been formed in situ making this structure a tidal tail from the SMC.  The SAGE-SMC (Surveying the Agents of Galaxy Evolution in the Tidally-Stripped, Low Metallicity Small Magellanic Cloud) program is a {\\em Spitzer} cycle 4 Legacy (285 hours, PI: K.\\ Gordon, PID: 40245) to map the full SMC (30~$\\sq\\arcdeg$) including the Body, Wing, and Tail using the Infrared Array Camera \\citep[IRAC,][]{Fazio04IRAC} and Multiband Imaging Photometer for {\\em Spitzer} \\citep[MIPS,][]{Rieke04MIPS} instruments on the {\\em Spitzer} Space Telescope \\citep{Werner04Spitzer}.  The SAGE-SMC program builds on the pathfinder S$^3$MC program \\citep{Bolatto07} that surveyed the inner $\\sim$3~$\\sq\\arcdeg$ of this galaxy.  The SAGE-SMC observations allow us to trace the life cycle of dust (and thereby gas) on a galaxy wide scale from their injection by late-type stars, through their sojourn in the violent ISM, until their involvement in the formation of stars. In addition, the infrared (IR) emission traces the global structure of the ISM on a galaxy-wide scale and the interrelationship of the various phases of the ISM. This survey provides a complete census of the massive star formation population in this low and spatially varying metallicity environment.  With much improved wavelength coverage, up to $\\sim$1000 times better point source sensitivity and $\\sim$13 times better angular resolution than the Midcourse Space Experiment ({\\em MSX}) and Infrared Astronomical Satellite ({\\em IRAS}) surveys, SAGE-SMC significantly improves our understanding of this important galaxy.  Combining the results from this SMC survey with the existing LMC \\citep[SAGE-LMC,][]{Meixner06} and Milky Way \\citep[GLIMPSE, MIPSGAL,][]{Churchwell09, Carey09} surveys provides a foundation for understanding the physics of the ISM, current star formation, and evolved-star mass loss as a function of metallicity.  This foundation is crucial for interpreting the observations of more distant galaxies like those in the SINGS \\citep{Kennicutt03SINGS}, SWIRE \\citep{Lonsdale03}, and GOODS \\citep{Dickinson03} {\\em Spitzer} Legacy programs.  The SAGE-SMC survey provides a crucial link in our understanding of galaxies during low metallicity, chemically-primitive stages.  This paper presents the overview of the SAGE-SMC Legacy program including the science motivation, observation details, data reduction, point source catalog creation, and publicly available data products.  In addition, results from all three areas (interstellar medium, star formation, and stellar mass loss) of scientific interest of the SAGE-SMC team are presented. ",
        "conclusions": "The motivation of the SAGE-SMC Sptizer Legacy was to study the lifecycle of the dust in the SMC; from birth in the winds of evolved stars, migration from the diffuse to dense ISM, and use as shielding and fuel in the formation of stars.  The SMC provides a unique laboratory for the study of the lifecycle of dust given its low metallicity and relative proximity.  The SAGE-SMC {\\em Spitzer} Legacy observed the full SMC from 3.6 to 160~\\micron\\ using the IRAC and MIPS instruments.  The region surveyed 30~$\\sq\\arcdeg$ encompassing the SMC Body (Bar and Wing) and Tail (high-density portion of the Magellanic Bridge).  The observations were carefully reduced and high quality mosaics and point source catalogs were produced.  These data products are available from the {\\em Spitzer} Science Center and IRSA Archive.  Initial results from the SAGE-SMC Legacy Team include finding that the gas-to-dust ratio in the SMC Tail region matches that of the SMC Body \\citep{Gordon09SMC}, probing the detailed CO and dust structure of an SMC molecular cloud to confirm that CO is significantly dissociated at cloud edges \\citep{Leroy09SMC}, studying the dust mass loss in the Galactic clusters NGC~362 \\citep{Boyer09} and 47~Tuc \\citep{Boyer10, McDonald11}, investigating the best monochromatic obscured SFR indicator \\citep{Lawton10}, studying the IR properties of known massive stars \\citep{Bonanos10}, and a detailed investigation of the star formation in NGC~602 \\citep{Carlson11}. A set of companion papers to this overview paper will investigate the interstellar dust (Bot et al. 2011, in prep.), evolved stars \\citep{Boyer11}, young stellar objects (Sewi\\'lo et al. 2011, in prep.), and HII/star formation regions (Lawton et al. 2011, in prep.).  The launch of the {\\em Herschel Space Observatory} has enabled the infrared dataset on the SMC (and LMC) to be extended in the submm through the HERITAGE Key Project \\citep{Meixner10}.  The future of Magellanic Clouds research in the infrared is very bright."
    },
    "1107/1107.1234.txt": {
        "abstract": "Studying star clusters offers significant advances in stellar astrophysics due to the combined power of having many stars with essentially the same distance, age, and initial composition.  This makes clusters excellent test benches for verification of stellar evolution theory.  To fully exploit this potential, it is vital that the star sample is uncontaminated by stars that are not members of the cluster. Techniques for determining cluster membership therefore play a key role in the investigation of clusters. We present results on three clusters in the \\kepler~field of view based on a newly established technique that uses asteroseismology to identify fore- or background stars in the field, which demonstrates advantages over classical methods such as kinematic and photometry measurements. Four previously identified seismic non-members in NGC~6819 are confirmed in this study, and three additional non-members are found -- two in NGC~6819 and one in NGC~6791. We further highlight which stars are, or might be, affected by blending, which needs to be taken into account when analysing these \\kepler~data. ",
        "introduction": "Determination of cluster membership is a crucial step in the analysis of stellar clusters.  Stars in an open cluster are thought to have formed from the same interstellar cloud of gas and dust, and hence share a common age and space velocity. % making them an ideal for verifying %stellar evolution theory. Cluster membership can therefore be inferred from the location of the stars along an isochrone in the color-magnitude diagram (photometric membership), and from their common space velocity (kinematic membership) measured as the line-of-sight radial velocity and the perpendicular proper motion. Recently, a new independent method was introduced by \\citet{Stello10} who performed an asteroseismic analysis based on the first month of data from the \\keplermission~\\citep{Koch10} to infer the cluster membership for a small sample of red giant stars in NGC~6819. Asteroseismology has the advantage that the oscillations in a star, which depend on the physical properties of the star's interior \\citep{Dalsgaard04}, are independent of stellar distance, interstellar extinction, and any random alignment between the space velocity of the cluster and field stars. In particular, the so-called average large frequency separation, \\dnu, between consecutive overtone oscillation modes depends on the mean density of the star, and the frequency of maximum oscillation power, \\numax, is related to its surface gravity and effective temperature. Both \\dnu~and \\numax~ are known to scale with %follow known scaling relations expressed by the basic stellar properties, $M$, $L$, and \\teff, \\citep{Ulrich86,KjeldsenBedding95} and can therefore be used to infer those properties without relying on detailed modelling of stellar interiors \\citep[e.g.~][]{Stello08,Kallinger10}.  We now have \\kepler~time series photometry that span 10 times longer than in the work by \\citet{Stello10}.  In this paper we are therefore able to present an asteroseismic membership analysis of a more comprehensive set of red giant stars in three open clusters within \\kepler's fixed field of view: NGC~6791, NGC~6819, and NGC~6811. We are further able to measure and make use of both \\dnu~and \\numax~for this purpose. In addition, to facilitate our inference on cluster membership and obtain more robust results we include an investigation of \\teff~found from different color indices and present a detailed analysis of blending.  Our membership results are compared with those from classical techniques.  We refer to our four companion papers for additional asteroseismic exploration of the same cluster data including (i) the determination of stellar mass and radius and cluster distances \\citep{Basu11}, (ii) verification of scaling relations for \\dnu~and \\numax~\\citep{Hekker10a}, (iii) derivation of a new scaling relation for oscillation amplitude \\citep{Stello11}, (iv) and mass loss properties of red giants during their transition between the hydrogen-shell and core-helium-burning phases \\citep{Miglio11}.  Like this paper, these studies are based on the global asteroseismic seismic properties, \\dnu, \\numax~and amplitude, while more detailed frequency analyses requires more data for these relative faint and crowded cluster stars. ",
        "conclusions": ""
    },
    "1107/1107.5229_arXiv.txt": {
        "abstract": "We continue our study of the application of the conformal gravity theory to galactic rotation curves. Previously we had studied a varied 111 spiral galaxy sample  consisting of high surface brightness galaxies, low surface brightness galaxies and dwarf galaxies. With no free parameters other than galactic mass to light ratios, we had found that the theory is able to account for the systematics that is observed in the entire set of galactic rotation curves without the need for any dark matter whatsoever. In the present paper we extend our study to incorporate a further 27 galaxies of which 25 are dwarf galaxies and provide updated studies of 3 additional galaxies that had been in the original sample, and again without dark matter find fully acceptable fits, save only for just a few galaxies that we find to be somewhat troublesome. Our current study brings to 138 the number of rotation curves of galaxies that have been accounted for by the conformal gravity theory. Since one of the primary ingredients in the theory is a universal contribution to galactic motions coming from matter exterior to the galaxies, and thus independent of them, our study reinforces one of the central concepts of the conformal gravity studies, namely  that invoking dark matter should be viewed as being nothing more than an attempt to describe global physics contributions in purely local galactic terms. ",
        "introduction": "\\label{s1}  As a possible alternative to standard Einstein gravity, Weyl introduced conformal gravity in the very early days of general relativity. It is an attractive theory in that it is a pure metric theory of gravity that possesses all of the general coordinate invariance and equivalence principle structure of standard gravity while augmenting it with an additional symmetry, local conformal invariance, in which  the action is left invariant under local conformal transformations on the metric of the form $g_{\\mu\\nu}(x)\\rightarrow e^{2\\alpha(x)}g_{\\mu\\nu}(x)$ with arbitrary local phase $\\alpha(x)$. Under such a symmetry the gravitational action is uniquely prescribed to be of the form (see e.g. \\cite{Mannheim2006}) \\begin{equation} I_{\\rm W}=-\\alpha_g\\int d^4x\\, (-g)^{1/2}C_{\\lambda\\mu\\nu\\kappa} C^{\\lambda\\mu\\nu\\kappa} \\equiv -2\\alpha_g\\int d^4x\\, (-g)^{1/2}\\left[R_{\\mu\\kappa}R^{\\mu\\kappa}-(1/3) (R^{\\alpha}_{\\phantom{\\alpha}\\alpha})^2\\right], \\label{182} \\end{equation} where \\begin{equation} C_{\\lambda\\mu\\nu\\kappa}= R_{\\lambda\\mu\\nu\\kappa} -\\frac{1}{2}\\left(g_{\\lambda\\nu}R_{\\mu\\kappa}- g_{\\lambda\\kappa}R_{\\mu\\nu}- g_{\\mu\\nu}R_{\\lambda\\kappa}+ g_{\\mu\\kappa}R_{\\lambda\\nu}\\right) +\\frac{1}{6}R^{\\alpha}_{\\phantom{\\alpha}\\alpha}\\left( g_{\\lambda\\nu}g_{\\mu\\kappa}- g_{\\lambda\\kappa}g_{\\mu\\nu}\\right) \\label{180} \\end{equation} is the conformal Weyl tensor and the gravitational coupling constant $\\alpha_g$ is dimensionless. With the conformal symmetry forbidding the presence of any $\\int d^4x\\, (-g)^{1/2}\\Lambda$ term in the action, the conformal theory has a control over the cosmological constant that the standard Einstein theory does not, and through this control one is able to both address and resolve the cosmological constant problem \\cite{Mannheim2009}. Similarly, with the coupling constant $\\alpha_g$ being dimensionless, unlike standard gravity conformal gravity is renormalizable; and with it having been shown \\cite{Bender2008,Mannheim2009} to be unitary at the quantum level, the theory is offered \\cite{Mannheim2011} as a consistent theory of quantum gravity in four spacetime dimensions.  With the conformal theory being a consistent, renormalizable quantum theory at the microscopic level, then just as with electrodynamics,  one is assured that its macroscopic classical predictions are reliable and will not be ruined by quantum corrections. Consequently, application of the theory to astrophysical phenomena allows one to test the theory. Early work in this direction was provided in \\cite{Mannheim1997} where the theory was used to fit the rotation curves of a set of 11 spiral galaxies, with the mass to light ratios ($M/L$) of the luminous optical disk of each galaxy being the only free parameters, and with no dark matter being required. More recently \\cite{Mannheim2010a,Mannheim2010b} a systematic, broad-based study of the rotation curves of a varied set of 111 galaxies (consisting of high surface brightness galaxies, low surface brightness galaxies and dwarf galaxies) was made (the set included the original 11 galaxies and updates where available), and again acceptable fitting was obtained with the optical disk mass to light ratios of each galaxy being the only free parameters, and again with no dark matter being required. In this paper we extend the study of \\cite{Mannheim2010a,Mannheim2010b} to a set of 30 galaxies that contains 20 of the dwarf galaxies described in \\cite{Swaters2009}, 2 low surface brightness galaxies (LSB), 5 of the dwarf galaxies described in \\cite{Oh2011}, and updates of 3 galaxies (2 dwarf galaxies and an LSB) that were in the 111 galaxy sample, and report findings of the same quality as before. With this latest study, the sample of galactic rotation curves  that can be accounted for by the conformal theory now runs to 138. ",
        "conclusions": ""
    },
    "1107/1107.0658.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract{We present the results of our 600 ks RGS observation as part of the multiwavelength campaign on Mrk 509. The very high quality of the spectrum allows us to investigate the ionized outflow with an unprecedented accuracy due to the long exposure and the use of the RGS multipointing mode. We detect multiple absorption lines from the interstellar medium and from the ionized absorber in Mrk 509. A number of emission components are also detected, including broad emission lines consistent with an origin in the broad line region, the narrow \\ion{O}{vii} forbidden emission line and also (narrow) radiative recombination continua. The ionized absorber consists of two velocity components ($v$ = $-$ 13 $\\pm$ 11 km s$^{-1}$ and $v$ = $-$ 319 $\\pm$ 14 km s$^{-1}$), which both are consistent with earlier results, including UV data. There is another tentative component outflowing at high velocity, $-$ 770 $\\pm$ 109 km $\\mathrm{s^{-1}}$, which is only seen in a few highly ionized absorption lines. The outflow shows discrete ionization components, spanning four orders of magnitude in ionization parameter. Due to the excellent statistics of our spectrum, we demonstrate for the first time that the outflow in Mrk 509 in the important range of $\\log\\xi$ between 1$-$3 cannot be described by a smooth, continuous absorption measure distribution, but instead shows two strong, discrete peaks. At the highest and lowest ionization parameters we cannot differentiate smooth and discrete components.} ",
        "introduction": "\\label{intro}  One of the main reasons to study active galactic nuclei (AGN) is to learn about feedback from the AGN to the galaxy and its direct environments. Feedback is a combination of enrichment (the spreading of elements into the interstellar and inter galactic media (ISM and IGM), momentum feedback (due to winds), and direct kinetic feedback (i.e. energy ejection into the ISM or IGM by jets). From recent observations on cooling clusters of galaxies \\citep[see e.g.][for an overview]{McNamara07}, as well as from recent insights into galaxy and AGN co-evolution \\citep{DiMatteo05,Elvis06,Bower09,Fabian10}, it has become clear that feedback from AGN is a crucial ingredient  the evolution of galaxies and clusters of galaxies. This is also seen in the so-called M-$\\sigma$ relation, which links the velocity dispersion of stars in the bulge to the mass of the Super-Massive Black Hole (SMBH) \\citep{Ferrarese00,Gebhardt00}. While we have a reasonable qualitative understanding of the feedback from relativistic jets \\citep[as observed in clusters of galaxies, see e.g.][]{Fabian03}, we still lack a quantitative picture of the feedback of the AGN on the galaxy and on its surroundings. \\\\ There is a broad ongoing effort to improve this, and recent work on broad absorption line (BAL) quasars shows that the mass outflow rates in these systems are 100s of solar masses per year and the kinetic luminosity involved is a few percent of the total bolometric luminosity \\citep{Moe09,Dunn10}. There have also been indications that some AGN harbor a highly ionized, massive, ultra-fast outflow, with velocities reaching up to 60\\,000 km s$^{-1}$ \\citep{Reeves03,Pounds09,Ponti09,Tombesi10a,Tombesi10b}. These outflows are hard to detect, however, and appear to be variable (because they are only present in some observations of a single source). These are extreme cases of outflows that are present in only a fraction of the total number of AGN. Whether feedback from less extreme outflows, such as those that are present in about 50 $\\%$ of the local Seyfert 1 galaxies is also important, remains an unsolved question.  If we can establish the impact that the outflow has on the galaxy in these local Seyfert 1 galaxies, this should allow us to extend the feedback estimates that we obtain to higher redshifts to the more powerful AGN, which we are unable to investigate with the current generation of X-ray grating spectrometers. However we first need to deal with the two main uncertainties concerning the outflows. These are the geometry of the inner region of an AGN and the location or origin of the outflow \\citep[see e.g.][]{Murray97,Krolik01,Gaskell07}. Earlier work has placed the outflow at various distances, and also the estimates for feedback can vary wildly \\citep[see e.g.][for some examples]{Behar03,Blustin05,Krongold07,Detmers08}. Therefore answering these two questions is the main goal of the Mrk 509 multiwavelength campaign.  Multiwavelength campaigns on AGN are crucial for gaining a complete understanding of the inner regions of these sources. Earlier multiwavelength campaigns focused mainly on abundance studies of the outflow \\citep[see Mrk 279,][]{Arav07} or on determining the outflow structure and location by combining UV and X-ray data e.g. NGC\\,5548 \\citep{Steenbrugge05}; NGC\\,3783 \\citep{Netzer03,Gabel03}. Our dedicated multiwavelength campaign on Mrk 509 is much more extensive than previous attempts. Our more intensive observations are ideal for locating the outflow, using the variability of the source and response of the ionized gas to determine its location \\citep[the use of variability to locate gas has been very successfully used in reverberation mapping of the BLR, see e.g.][for an overview of the method and for the latest results]{Peterson00,Denney10}.  Apart from the location and kinematics, one of the other important questions regarding the outflow is what the ionization structure is. Earlier studies have reported different results. The outflow in NGC 5548 appears to be a continuous distribution of column density vs. log $\\xi$ \\citep{Steenbrugge05}. NGC 3783, on the other hand, shows different separate ionization components, all in pressure equilibrium \\citep{Krongold03}. In Mrk 279 the situation appears to be more complex, because a nonmonotonous, continuous distribution provides the best description to the data \\citep{Costantini07a}. Recently, \\citet{Holczer07} and \\citet{Behar09} have shown that for most local Seyfert 1 galaxies with an outflow, a continuous distribution of column density vs. ionization parameter is the best description of the data. What they also show is that there are distinct $\\xi$ values where no is gas present. They interpret these gaps as thermal instabilities that cause the gas to rapidly cool or heat and then shift to other ionization parameters. What is clear from these studies is that there is no single model that describes all the observed outflows. High$-$quality, high$-$resolution spectra of the outflows are crucial for investigating the structure, since it can be the case (as in Mrk 279) that some components of the outflow have very low column density, which would otherwise escape detection.  Mrk 509 is one of the best studied local AGN, and due to its large luminosity ($L(1-1000\\,\\mathrm{Ryd})$ = 3.2 $\\times$ 10$^{38}$ W), it is also considered to be one of the closest QSO/Seyfert 1 hybrids. Earlier work on the outflow in the X-ray regime has revealed that it consists of a wide range of ionization components, but lacks the very high and also very low ionized gas \\citep[weak Fe UTA and no \\ion{Si}{xiv} Ly$\\alpha$,][]{Yaqoob03}. The outflow has been described using three ionization components, each with a different outflow velocity \\citep{Smith07}; however, the exact outflow velocities differ between different publications, most likely due to a limited signal$-$to$-$noise ratio. \\citet{Detmers10} have analyzed three archival observations of Mrk 509 with \\textit{XMM-Newton}. They also find three components for the outflow, although with different velocities than \\citet{Smith07}. Including the EPIC-pn data and improving the relative calibration between RGS and EPIC-pn achieved increased sensitivity. With this improvement they were able to detect variability in the highest ionization component, constraining the distance of that component to within 0.5 pc of the central source. Another point of interest is that there have been indications of an ultra-high velocity outflow as seen through the \\ion{Fe}{K} line \\citep{Cappi09,Ponti09}. This outflow could make a potentially large contribution to feedback, as the velocity is very high, although it appears to be transient \\citep{Ponti09}.  This work is the third in a series of papers analyzing the very deep and broad multiwavelength campaign on Mrk 509. The complete campaign details are presented in \\citet{Kaastra2011a}, hereafter Paper I. Here we present the main results obtained from the stacked 600 ks \\textit{XMM-Newton} RGS spectrum. With this spectrum, we are able to characterize the properties of the ionized outflow in great detail (velocities, ionization states, column densities, density profile, etc.). Other features detected in the spectrum ( emission lines, Galactic absorption, etc.) will not be discussed in detail here. Because different physics are involved, we will discuss them in future papers in this series.  This article is organized as follows. Section \\ref{model} briefly describes the data reduction for obtaining the stacked spectrum, and we show the spectral models that we use to describe the data. The spectral analysis and the results are presented in Sect. \\ref{spectral}. We discuss our results in Sect. \\ref{discussion} and present our conclusions in Sect. \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have presented one of the highest signal-to-noise RGS spectra of an AGN. With the almost unprecedented detail in this dataset, we could detect multiple absorption systems. The ionized absorber of Mrk 509 shows three velocity components, one at $-$13 $\\pm$ 11 km s$^{-1}$, one at $-$319 km s$^{-1}$, and a tentative high$-$velocity component at $-$ 770 km s$^{-1}$. The first two components are consistent with the main absorption troughs in the UV. Thanks to the high$-$quality spectrum and the accurate column densities obtained for all ions, for the first time it has been shown clearly that the outflow in Mrk 509 in the important range of $\\log\\xi$ between 1$-$3 cannot be described by a smooth, continuous absorption measure distribution, but instead shows two strong, discrete peaks. At the highest and lowest ionization parameters, we cannot distinguish between smooth and discrete components. We also have found indications of an increasing outflow velocity versus ionization parameter. Large, dedicated multiwavelength campaigns such as this are the way forward, as this is currently the best method to investigate and characterize the outflows in the local Seyfert galaxies."
    },
    "1107/1107.0583_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  Young star clusters (YSCs) are considered a powerful tool to study the star formation history of their hosts.  However, many assumptions and constraints on the evolutionary properties are needed in order to compare YSCs properties and the bulk of the host stellar population. In nearby galaxies, it is possible to study both the resolved stellar population and cluster population \\citep[e.g.][]{2011A&A...529A..25S}. In farther galaxies, however, the study of YSCs is conducted with an analysis of their integrated light at different wavelengths under the assumption of  an instantaneous burst, i.e. single stellar population.  The luminous (M$_B< -17.0$, corresponding to total stellar masses of $\\sim 10^{9-10}$ M$_{\\odot}$) blue compact galaxies (BCGs) show clear signatures of interactions and/or mergers, and the numerous observed massive YSCs are likely the result of these encounters \\citep{2003A&A...408..887O, 2010MNRAS.407..870A, 2011MNRAS.414.1793A, 2011MNRAS.tmp..739A}. The very bright ultraviolet and optical luminosities of these systems suggest rather low dust and metallicity content. Spectra dominated by emission lines clearly demonstrate that BCGs are undergoing a starburst episode. The youth of the burst episode is also observed in the recovered age distribution of the star clusters, which shows a peak of cluster formation younger than 5 Myr.  The analysis of the young cluster populations in BCGs is quite challenging due to the rapid evolution a cluster experiences during the first 10 Myr (still partly in an embedded phase). Moreover, this analysis is based on the integrated luminosities of the clusters, which appear unresolved at the distance of the targets.  Observations of resolved newly born star clusters in the Milky Way and in the Magellanic Clouds reveal that these are quite complex systems. A cluster  forms in a hierarchical medium, from the fragmentation and collapse of giant molecular cloud complexes \\citep{2010IAUS..266....3E}. This implies that stars form in a fractal distribution and only in some cloud cores are their number densities is high enough to eventually form a bound cluster \\citep{2010MNRAS.409L..54B}. This picture is confirmed also by dynamical studies. \\citet{2011MNRAS.410L...6G} observed a continuous distribution between associations (unbound) and clusters (bound) at very young ages, and a clear distinction at older stages, when the crossing time of unbound systems exceeds their stellar age.  In a newly born cluster, the massive and short-lived stars, rapidly reach the main sequence and produce strong winds and UV radiation, which ionizes the intracluster gas and create bubbles and shells. These H{\\sc ii} regions surround the optically bright core of stars and significantly contribute to the integrated fluxes. However, a large fraction of the stars is still accreting material from their dusty disks (young stellar objects, YSOs) or contracting (in the pre-main sequence phase, PMS). Moreover, the edges of the clusters are places for triggered \\citep{2011arXiv1101.3112E} and progressive star formation \\citep[e.g.][]{2007ApJ...665L.109C}, which can explain the observed age spread in the  star forming regions \\citep[e.g.][]{2010ApJ...713..883B}. ",
        "conclusions": ""
    },
    "1107/1107.2779_arXiv.txt": {
        "abstract": "The vector curvaton paradigm is reviewed. The mechanism allows a massive vector boson field to contribute to or even generate the curvature perturbation in the Universe. Contribution of vector bosons is likely to generate statistical anisotropy in the spectrum and bispectrum of the curvature perturbation, which will soon be probed observationally. Two specific models for the generation of superhorizon spectra for the components of an Abelian vector field are analysed. Emphasis is put on the observational signatures of the models when the vector fields play the role of vector curvatons. If future observations support the vector curvaton mechanism this will open a window into the gauge field content of theories beyond the standard model. ",
        "introduction": "Cosmic inflation is arguably the most compelling way to overcome, or at least ameliorate, the so-called horizon and flatness problems of the hot big bang cosmology. However, these problems are successfully addressed by all models of inflation, provided that the inflationary expansion lasts long enough. Therefore, discrimination between inflation models is based on another important consequence of inflationary expansion, namely the generation of the curvature perturbation $\\zeta$ in the Universe, which is responsible for the formation of structures such as galaxies and galactic clusters and which is reflected onto the Cosmic Microwave Background (CMB) radiation through the Sachs-Wolfe effect.  The latest CMB observations appear to confirm the ``vanilla'' predictions of inflation with respect to $\\zeta$: scale-invariance, Gaussianity and statistical homogeneity and isotropy. However, a period of accelerated expansion of space (this is the definition of cosmic inflation) is not enough to guarantee scale-invariance for the curvature perturbation. Indeed, inflation is required to be of the quasi-de Sitter type, where the density of the Universe remains roughly constant. Also, the CMB observations seem to suggest that exact scale-invariance is not favoured (although it is not ruled out) with the spectral index of $\\zeta$, satisfying \\mbox{$n_s-1=-0.037\\pm0.014$} (at 1$\\sigma$) \\cite{wmap7} when $\\Lambda$CDM cosmology is assumed.% \\footnote{Exact scale invariance corresponds to \\mbox{$n_s=1$}.} Most likely, this reveals some dynamics for inflation (the density is not exactly constant), which is indeed expected by model-builders. Similarly, a high degree of Gaussianity in the curvature perturbation is expected since this reflects the randomness of quantum fluctuations, on which the particle production process is feeding, in order to generate the superhorizon spectrum of the fields which eventually give rise to $\\zeta$. However, the Gaussianity of $\\zeta$ crucially depends on the linearity of the process which translates these field perturbations to $\\zeta$. This is quantified through the so-called non-linearity parameter $f_{\\rm NL}$. The latest CMB observations provide a hint of non-zero non-Gaussianity since, in the squeezed configuration, \\mbox{$f_{\\rm NL}=32\\pm21$} (at 1$\\sigma$) \\cite{wmap7}, which again appears to deviate from the ``vanilla'' prediction.\\footnote{Note that $\\zeta$ is indeed predominantly Gaussian because the bispectrum $B_\\zeta$ is related to the power spectrum $P_\\zeta$ roughly as \\mbox{$B_\\zeta\\sim f_{\\rm NL}P_\\zeta^2$} (see Eq.~(\\ref{fNLdef})), with the observations suggesting \\mbox{$\\calp_\\zeta=4.8\\times 10^{-5}$} for the spectrum \\cite{wmap7}.}  In the same spirit, observations suggest that there may be deviations from statistical homogeneity, since there seem to be a difference in the power of $\\zeta$ as large as 10\\% between hemispheres in the CMB \\cite{stinhom}. Finally, statistical isotropy in $\\zeta$ is also questioned by observations. Indeed, there is tantalising evidence of a preferred direction on the microwave sky. This is the so-called ``Axis of Evil'' observation \\cite{AoE}, which amounts to an alignment of the quadrupole and octupole moments in the CMB, which is statistically extremely unlikely \\cite{anistat} and has been shown to persist beyond foreground removal \\cite{forg}.  Thus, we see that the precision of the cosmological observations is such that begins to enable us to explore beyond the ``vanilla'' predictions of inflation and use the observed deviations from them to discriminate between classes of models and paradigms. There is already a huge literature on deviations from exact scale-invariance and Gaussianity for the curvature perturbation. In this paper we discuss possible deviations from statistical isotropy and we present a compelling paradigm for their generation; namely the Vector Curvaton Paradigm.  Throughout the paper we consider a metric with negative signature and use natural units where \\mbox{$c=\\hbar=k_B=1$} and Newton's gravitational constant is \\mbox{$8\\pi G=m_P^{-2}$}, with \\mbox{$m_P=2.4\\times 10^{18}\\,$GeV} being the reduced Planck mass. ",
        "conclusions": "The high precision cosmological observations enable cosmologists to investigate beyond the ``vanilla'' predictions of inflation and thereby discriminate between inflation models. A new such observable is statistical anisotropy, which amounts to direction dependent patterns in the CMB (and possibly large scale structure too). Statistical anisotropy is within the reach of the forthcoming observations of the Planck satellite mission, which is expected to release its first cosmological data in the beginning of 2013. Currently, the latest CMB observations allow statistical anisotropy in the spectrum as much as 30\\% (\\mbox{$g\\lsim 0.3$}). Planck will reduce this bound down to 2\\% if statistical anisotropy is not observed. It should be noted here that, even if the spectrum is weakly statistically anisotropic, the bispectrum can be predominantly anisotropic with \\mbox{${\\cal G}\\gg 1$}. This means that, if non-Gaussianity is indeed observed without a strong angular modulation of $f_{\\rm NL}$ on the microwave sky, then all models which predict \\mbox{${\\cal G}\\gg 1$} will be ruled out.  Vector boson fields are natural candidates for the generation of statistical anisotropy in the curvature perturbation, because they are expected to pick a preferred direction when homogenised by inflation. The vector curvaton paradigm offers a simple, generic and effective mechanism for the direct contribution of vector boson fields to the curvature perturbation. The mechanism assumes a Proca vector field, whose zero-mode begins oscillating when the field becomes heavy, after the end of inflation. As shown, the oscillating zero-mode behaves as a pressureless isotropic fluid and can (nearly) dominate the Universe without generating an excessive anisotropic stress. When doing so it imposes its own contribution to the curvature perturbation, in accordance to the curvaton mechanism. We should point out here that this perturbation is scalar in nature, because it is due to the perturbed value of the density of the vector field, which is a scalar quantity. A considerable advantage of the vector curvaton mechanism is that it does not rely on an interaction of any kind between the vector field and the inflaton sectors. This allows the vector field to correspond to physics at a much lower energy scale (e.g. TeV physics) than the scale of inflation, which may connect with observations in collider experiments such as the LHC.  The particle production process, through which the vector curvaton obtains a superhorizon spectrum of perturbations, is in general anisotropic. This is because the vector boson field has more than one degrees of freedom (three if massive), for which the efficiency of the particle production is in general different. Thus, the curvature perturbation contributed by a vector curvaton is, in general, statistically anisotropic. If particle production is indeed isotropic (at least at the level allowed by the observations) the vector curvaton mechanism can generate the curvature perturbation in the Universe from vector fields alone without directly involving any fundamental scalar fields. If, however, statistical anisotropy is indeed observed, then we {\\em have} to go beyond scalar fields to explain the observations. This means that, the observation of statistical anisotropy in the CMB, may probe the gauge field content of theories beyond the standard model.  There are some related issues which we did not go into in this paper. For example, studies of the trispectrum in vector field scenarios \\cite{Ytri}, or of one-loop contributions \\cite{Y1loop}. Note also, the possibility that the vector curvaton is non-Abelian is investigated in Ref.~\\cite{nonAbelian}, while the contribution to the curvature perturbation by P-forms can be found in Ref.~\\cite{Pforms}.  The interest now is in finding realistic candidates in theories beyond the standard model, which can play the role of the vector curvaton. Examples can include the supermassive gauge bosons of grand unified theories \\cite{WDL}\\footnote{These can be associated with the simultaneous generation of a primordial magnetic field \\cite{pmf}.} or the vector fluxes on probe branes in the context of DBI-inflation \\cite{DBIvc}. Another promising idea is investigating the possibility that the vector bosons associated with gauged axions can play the role of the vector curvaton. In this case, the generation of parity-violating statistical anisotropy is possible \\cite{anomaly}."
    },
    "1107/1107.1857_arXiv.txt": {
        "abstract": "{Sun-grazing comets almost never re-emerge, but their sublimative destruction near the sun has only recently been observed directly, while chromospheric impacts have not yet been seen, nor impact theory developed.} {We seek simple analytic models of comet destruction processes near the sun, to enable estimation of observable signature dependence on original incident mass $M_o$ and perihelion distance $q$.} {Simple analytic solutions are found for $M(r)$ versus $q$ and distance $r$ for insolation sublimation and, for the first time, for impact ablation and explosion.} {Sun-grazers are found to fall into three ($M_o,q$) regimes: sublimation-, ablation-, and explosion-dominated. Most sun-grazers have $M_o$ too small ($< 10^{11}$ g) or $q$ too large ($>1.01R_\\odot$) to reach atmospheric densities ($n>2.5\\times 10^{11}$/cm$^3$) where ablation exceeds sublimation. Our analytic results for sublimation are similar to numerical models. For $q<1.01R_\\odot, M_o>10^{11}$ g, ablation initially dominates but results are sensitive to nucleus strength $P_c=10^6P_6$ dyne/cm$^2$ and entry angle $\\phi$ to the vertical.  Nuclei with $M_o \\preceq 10^{10}(P_6\\sec\\phi)^3$ g are fully ablated before exploding, though the hot wake itself explodes. For most sun-impactors $\\sec\\phi \\gg 1$ (since $q\\sim r_*$), so for $q$ very close to $r_*$ the ablation regime applies to moderate $M_o \\sim 10^{13-16}P_6^3$ g impactors unless $P_6\\preceq 0.1$. For higher masses, or smaller $q$, nuclei reach densities $n> 2.5\\times 10^{14}P_6$/cm$^3$ where ram pressure causes catastrophic explosion.} {Analytic descriptions define ($Mo,q$) regimes where sublimation, ablation and explosion dominate sun-grazer/-impactor destruction. For $q\\prec 1.01R_\\odot, M_o\\succeq 10^{11}$ g nuclei are destroyed by ablation or explosion (depending on $M_o\\cos^3\\phi/P_c$) in the chromosphere, producing flare-like events with cometary abundance spectra. For all plausible $M_o,q$ and physical parameters, nuclei are destroyed above the photosphere.} ",
        "introduction": "\\label{Introduction}  Close sun-grazing comets are discovered almost daily by white light coronographs (e.g. SoHO LASCO). Most are small and fully sublimated by insolation at a few solar radii $R_\\odot$ while almost none have re-emerged (Marsden 2005). The majority have perihelion distances $q$ well above $R_\\odot$ but some have $q\\approx$ or $\\prec R_\\odot$ (e.g. Marsden 2005 and http://www.minorplanetcenter.org/mpec/RecentMPECs.html.) The death of a comet at $r\\sim R_\\odot$ has been seen directly only very recently (Schrijver et al 2011) using the SDO AIA XUV instrument. This recorded sublimative destruction of Comet C/2011 N3 as it crossed the solar disk very near perihelion $q=1.139R_\\odot$. The next challenge in studying the demise of close sun-grazers will be to catch one of the even rarer cases of chromospheric impact ($q<1.01R_\\odot$, $M_o>10^{11}$ g). As anticipated by Weissman (1983) and shown below, these undergo explosive destruction in the dense chromosphere. Understanding the destruction processes, and their radiation signatures, are essential steps in searching for and modeling these.  The processes leading to sublimation of the icy conglomerate mix (Whipple 1950) of cometary nuclei, and in some cases their splitting and fragmentation, were considered by  Huebner (1967), Weissman and Kieffer (1981), Weissman (1983), Iseli et al. (2002), Sekanina (2003) and others. These models essentially solve for the insolative sublimation mass loss rates of icy-conglomerate mixes (the dust being carried away in the flow of these evolved components). They variously allow for the complicating factors of rotation, albedo, insulating surface dust layers, radiative cooling, interior thermal conduction, and fragmentation by tidal, thermal and volatile explosion effects. Huebner (1967) and Iseli et al. (2002), for example, found that, for high sublimation rates near the sun, these effects were secondary, the mass loss being reasonably approximated by a pure sublimation description : mass loss rate = heating rate/latent heat. Using SoHO data, Sekanina (2003) addressed in detail how mass loss rates and fragmentation relate to cometary light curves via atomic line emission (e.g. by sodium) and by dust scattering of sunlight, though emphasising that most of the mass remains in a primary fragment. None of these studies considered ablation or ram-pressure driven explosion due to solar atmospheric impact (which we show below are negligible till $r\\preceq1.01R_\\odot$) though  Weissman (1983) had remarked {\\it \"ultimate destruction of the nucleus [of sun-impacting comets] likely results from the shock of hitting the denser regions of the solar atmosphere\"}. It is well known in the planetary physics community that ablation and explosion are central processes in comet-planetary atmosphere impacts (e.g. Carslon et al. 1997)  Here we revisit the theory of sun-grazer sublimation then develop the first estimates of the much higher rates of mass loss by ablation for the rarer cases which reach  $\\preceq1.01R_\\odot$. This regime (Sections 3.4, 6.1, 6.2) resembles that of comet-planetary impacts though with some differences. One is that almost all sun-grazing comets belong to secondary comet groups (mostly Kreutz) formed from primordial comets. A significant number of the latter must have $q<R_\\odot$ (Hughes 2001) and large enough mass ($M_o\\succeq 10^{11}$ g - see Section 5, Eqn. (\\ref{Mominq})) to survive sublimation down to the intense ablation/explosion regime. However only a small fraction of group comets (Biesecker et al. 2003, Knight et al 2010) come that close or are that massive, larger sun-grazers mostly having $q\\succeq 1.5 R_\\odot$. Comet C/2010 E6 (STEREO) discussed by Raftery et al. (2010) came close,  having $q\\approx 1.02 R_\\odot$ while the destruction of Comet C/2011 N3 seen by SDO (Schrijver et al 2011) on the solar disk was solely by sublimation as it had $q\\approx 1.14R_\\odot$. Most of the few group comets reaching the ablation/explosion regimes will have small mass and very shallow incidence angles (since $q\\approx R_\\odot$) with implications for the relative importance of ablation and explosion. By analogy with models of planetary atmospheric impacts (e.g. Carlson et al. 1997) , chromospheric impactors are expected to lose mass, momentum and kinetic energy very rapidly, the soaring temperature and pressure leading to destructive detonation of the nucleus and its wake. This 'air burst' should be followed by explosive expansion of the resulting 'fireball' of cometary and solar atmospheric material. The total masses and energies of these explosive chromospheric events are in roughly the same range as those of solar magnetic flares so they could be termed 'cometary flares'.  Section 2 discusses relevant nucleus parameters values while Section 3 estimates the relative importance and time/distance-scales of nucleus mass loss, explosion, and deceleration processes. We then consider in some detail approximate analytic treatment of mass loss (Section 4) purely by sublimation (Section 5) and by ablation and explosion (Section 6). ",
        "conclusions": "\\subsection{Summary of Nucleus Destruction Regimes} We have shown there to be three regimes of nucleus destruction, depending on the values of $M_o,q$ for given values of $\\tilde{\\rho}, \\tilde{{\\cal L}}$ (for which we use here $\\tilde{\\rho}=0.5, \\tilde{{\\cal L}}=1$), namely \\begin{itemize} \\item {\\bf (a) Sublimation dominated} (Section 5.2.2, Eqn. (\\ref{Mominq}))  Nuclei of $p \\succeq x_*, M_o \\preceq M_{**}(p)$ given by \\begin{equation} \\label{M1ofp} M_1(p) =M_{ominq}^{rad}=\\frac{M_{*}}{p^{3/2}}\\approx \\frac{10^{11}}{p^{3/2}}~{\\rm g} \\end{equation} are completely sublimated before they reach perihelion, (or $r_*$). This is by far the commonest regime for group comets.  \\item {\\bf (b) Ablation dominated} (Sections 3.4.2, 6.1)  Since $n=n_{**}=2.5\\times 10^{14}P_6$ at $x_{**}$, where atmospheric ram pressure starts to exceeds nucleus strength, nuclei with $x_{**}<p<x_*$ and $M_1(p)<M_o<M_2(p)$ (i.e. of sufficiently low mass or shallow entry angle) lose all their mass to ablation before explosion with  \\begin{eqnarray} \\label{M2ofp} &&M_2(p) = M_{abl}^{max}(p)=\\nonumber \\\\ &&M_{**}P_c^3(1-p/x*)^{-3/2}\\approx 10^{10}P_6^3(1-p/x*)^{-3/2}{\\rm g} \\end{eqnarray} Since most larger sun-grazers have $p>x_*$, these are rarer than comets in the sublimation regime.  \\item {\\bf (c) Explosion dominated} (Section 6.2)  Nuclei with $p<x_{**}$ and $M_o>M_2(p)$ (ie of sufficiently high mass or steep entry) undergo ram pressure driven explosion. These are also rare compared to sublimative destruction. \\end{itemize} These domains in the plane of $(p/x_*,M_o)$  are summarized Figure 5 with the $M_o$ axis scaled with respect $P_6^3$. $P_6=1$ represents relatively strongly bound nuclei. For $P_c\\sim 10^5$ dyne/cm$^2$ the ablation regime becomes confined to about $M_o\\succeq 10^{11} {\\rm g}, 0.99\\preceq p/x_*\\preceq 1$. Though this is very narrow in $p$ space, how common ablation domination is compared to explosion domination depends on the distribution of $p/x_{**}$ and is very sensitive to the value of $P_c$. For very weakly bound bodies ($P_6 \\ll 0.1$) ablation would only dominate briefly prior to explosion totally dominating. Solid stony or iron bodies with $P_6\\succeq 10$ would, unless extremely massive, be wholly in the ablative destruction regime.   \\begin{figure} \\resizebox{\\hsize}{!} {\\includegraphics{figure5.eps}} \\caption{Domains of destruction of comet nuclei in the plane ($M_o/P_c^3,p/x_*=\\sin \\phi_*$) for entry angle $\\phi_*=\\arcsin(p/x_*)$ to the vertical at $r=r_*$. NOTE the change from linear to log scale on the $p/x_*$ axis across the point $p/x_*=1$ so as to display the very rapid variation in the range $p/x_*\\le 1$ The axes are scaled with respect to $P_6$ as discussed in section 2.2. For $P_6=0.1$ the ablation regime becomes a slender strip at $0.99\\preceq p/x_*\\preceq1$} \\label{fig5} \\end{figure} \\subsection{Visibility of Sun-grazer destruction by sublimation} It was shown in Section 3.3 that destruction by sublimation of low mass nuclei even very near the sun occurs over scales $\\sim R_\\odot$ in space and $\\sim 10^{3-4}$ s in time, making them hard to observe against solar atmospheric emission and scattering. However, the abrupt collisional stopping of sublimated material in the wake could raise its initial temperature by up to around $10^7$ K which may make it visible in the XUV by charge exchange or thermal line emission if it does not cool too quickly by radiation, conduction or expansion and if its emission measure $\\int_Vn^2dV$ is high enough. (The first observation, in XUV, of such an event was made by Schrijver et al., 2011, using SDO AIA during the revision stage of this paper). A $10^{12}$ g mass totally sublimated over a $0.5 R_\\odot$ path into a conical wake of half angle say $10^{-2}$ radians would have an emission measure of order $3\\times 10^{44}$ cm$^{-3}$ and a mean density of $\\sim 5\\times 10^8$ cm$^{-3}$ in a solar plasma environment of roughly similar density.  \\subsection{Signatures of Sun-impactor Destruction} \\subsubsection{Mass, Energy and Temperature of the Fireball} The kinetic energy initially released in the lower solar atmosphere by sun impacting cometary nuclei with $M_o=10^{11}-10^{19}$ g  is $5\\times (10^{26}-10^{34})$ erg, similar to the range of the magnetic energy release from 'nano'-flares to 'super'-flares. If 50\\% of ${\\cal E}_{kin}$ went into heating the comet material the temperature attainable would be $T>m_pv_\\odot^2/4k\\approx 10^7 K$ where $k$ is Boltzmann's constant. This is of order the solar escape temperature since it is generated by infall of matter effectively from $\\infty$.  Flare masses ejected are of the order of $10^{12-16}$ g,  also comparable with those of larger impacting cometary nuclei. Thus a comet impact in the dense chromosphere should produce phenomena somewhat resembling a solar (magnetic) flare, though the initial comet volume is much smaller, its density much higher, and the heating rise time faster ($\\preceq 10$ s). Because of the initially small size scale and high density and pressure we might expect rapid initial cooling (seconds) by radiation, expansion and conduction, followed by a radiative/conductive decay phase more like that of solar flares ($\\sim 10^{2-3}$ s) as the fireball expands outward and upward, sweeping solar plasma with it (cf. the Carlson et al. 1997 study of the Jupiter impact fireball).  \\subsubsection{Cometary Flare Size, Emission Measure and Optical Depth}  If optically thin, the radiative output of a hot plasma of volume $V$ depends on its temperature and its emission measure $EM= \\int_Vn^2dV\\approx a^3(M_o/\\mu_c m_p)^2\\approx 10^{48} (M_o/10^{12})^2/(a/10^8)^3/\\mu_c$ cm$^{-3}$ for a nucleus of mass $M_o$ when it has expanded to size $a$. (Here $\\mu_c$ is the mean comet mass per particle in units of $m_p$). Explosion of a nucleus occurs over a few scale heights $\\approx 10^8$ cm vertically, traversed in $\\sim 10$ s at speed $v_\\odot\\cos \\phi$. If the lateral expansion speed of the hot wake were similar then the lateral dimension of the exploding wake would also be a few scale heights. Its optical depth for a process with absorption cross-section $\\sigma$ would then be $\\tau \\approx \\sigma M_o/m_pa^2 = (M_o/10^{12})(\\sigma/10^{-20})/\\mu_c/(a/10^8)^2$. Thus when a nucleus of, say, $M_o=10^{12}$ g ($a_o\\sim 10^4$ cm) becomes optically thin ($a=10^6$ cm) for $\\sigma=10^{-20}$ cm$^2$, it would have an $EM\\approx 10^{48}$ cm$^{-3}$ for that emission. This $EM$ is only for the cometary material itself, but the $EM$ of the heated atmosphere may also be important since an atmospheric mass $\\succeq M_o$ is involved in decelerating it. Since $H\\sim$1 arcsec we may expect to see initially a source of a few arcsec expanding at about 0.1 arcsec per sec at chromospheric altitudes and atmospheric density $\\simeq 10^{13-16}$ cm$^{-3}$ depending on $M_o$ as per Figure 4. XUV line spectra should exhibit highly non-solar abundances (eg O:H ratio) as should any matter ejected into space (cf. Iseli et al 2002). There may also be a variety of nonthermal radio and other signatures (e.g. charge exchange - e.g. Lisse et al. 1996 - and impact polarized spectrum lines - e.g. Fletcher and Brown 1992, 1995) arising from plasma phenomena driven by the highly supersonic expansion, such as shock and/or turbulent acceleration of electrons and ions and charge separation currents as the ablated matter decelerates in the atmosphere - cf. review by P\\'{e}gauri\\'{e}r (2007) of fusion pellet ablation physics.  \\subsubsection{Cometary Sunquakes and other solar disturbances}  The phenomenon of flare induced sunquakes - waves in the photosphere -  discovered by Kosovichev and  Zharkova (1998) and now widely studied (e.g. Kosovichev 2006) should also result from the momentum impulse delivered by a cometary impact. All such impacts, however small the comet mass, involve a huge kinetic energy density (ram pressure) $\\rho v_\\odot^2/2\\approx 10^{15}$ erg/cm$^3$. This is $\\sim 10^{10}$ times the thermal energy density (pressure) of the photosphere or the magnetic energy density in a sunspot, being the energy density of a 40 MG magnetic field! Even when the kinetic energy is converted by ablation and ram pressure to heat and kinetic energy of explosion, it is initially spread over only a few scale heights. Hence the explosive airburst energy density of a comet like a Shoemaker-Levy 9 ($10^{15}$ g) is about $10^6$ erg/cm$^3$ equivalent to that of a 5 kG field and still more than the thermal energy density of the photosphere. Thus, while the total energy of most impacting sun-grazers is small compared to that of large flares and CMEs their energy density is so high that local disruption of magnetic fields and triggering of larger scale events are not impossible. \\subsection{Conclusions} Our simple analytic treatment of comet nucleus sublimation gives results for mass loss in reasonable agreement with previous numerical simulations . We have proved that nuclei reaching $p=q/R_\\odot \\precsim r_*/R_\\odot\\approx 1.01$ undergo intense bombardment by solar atmospheric ions (energy 2 keV per nucleon) the rate increasing exponentially with depth on a scale of around 500 km. It follows that, above depth $r_*$ where $n=n_*=2.5\\times 10^{11}$/cm$^3$, evaporative destruction is essentially by sublimation only. On the other hand, the behavior of any nucleus surviving sublimation to below this height is dominated by atmospheric collisional effects - ablation and ram pressure driven explosion. This creates an exploding air burst followed by a fireball spreading and rising through the atmosphere similarly to the Shoemaker Levy 9 impacts on Jupiter (e.g. Carlson et al. 1997).  Because of the exponential distribution of atmospheric density and the small atmospheric scale height $H$ (compared to $R_\\odot$) the terminal heights of these cometary flares is only logarithmically sensitive to the incoming cometary mass (and to the density and latent heat of the nucleus). It only ranges from the upper chromosphere for small to moderate comet masses with typically shallow entry angles to the low chromosphere for very massive steep entry comets.  The ablated matter (atoms and small particles) decelerates very quickly in collisions with the atmosphere, depositing its kinetic energy and momentum over a few scale heights in tens of seconds, heating the debris and local solar plasma to X-ray emitting temperatures similar to solar flares and with comparable plasma masses and emission measures. Observing such very small impulsive but very luminous 'cometary flares' poses a fascinating challenge but will enable new diagnostics of cometary element abundances from the highly non-solar flash emission spectra. In addition the downward momentum impulse should generate cometary sunquakes as observable ripples in the photosphere akin to those found in conventional magnetically-energized solar flares. In the case of impacts by the most massive comets ($10^{20}$ g or so) the cometary flare energy release ($2\\times 10^{35}$ erg) is much larger than that of the largest solar flares ever observed. An impact of this magnitude would have very significant terrestrial effects."
    },
    "1107/1107.3580_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro} Compact galaxy groups (CGs) occupy an interesting part of the parameter space that pertains to the clumping of matter in the universe. On the one hand, they lie at the low end tail of the distribution of membership size, as they contain few galaxies -- typically three or four \\citep{hickson82}. On the other hand, their compactness places them on the high end of the number density distribution, similar to that in the centres of galaxy clusters \\citep{dressler80,brad90}.  In addition, CGs display low velocity dispersions of $\\sigma_\\textup{\\scriptsize CG}\\sim250~$km$\\,$s$^{1}$ \\citep{tago08,asqu}, compared to galaxy cluster dispersions of $\\sigma_\\textup{\\scriptsize cluster}\\sim750~$km$\\,$s$^{1}$ \\citep{binggeli87,the86,asqu}. This dynamical situation lengthens interactions, when such events occur, or forces galaxies to a state of quasi-secular evolution: while they may not interact physically with their neighbours, they are always affected by them dynamically \\citep{martig08,isk10}. In terms of galaxy evolution, they present potential precursors of isolated ellipticals. Furthermore, they allow for close studies of galaxy evolution and morphological transformation, given their relatively few degrees of freedom (as compared with galaxy clusters and their hundreds of members).  \\begin{figure}[!t] \\center \\includegraphics[width=0.8\\textwidth, angle=270]{Konstantopoulos_I_S_Fig1.eps} \\caption{ The evolutionary diagram proposed by \\citet{isk10}. Gas content decreases from left to right, according to the definition of \\citet{johnson07}. The top sequence includes groups that hold the H\\,{\\sc i} gas within the member galaxies, while the CGs of the bottom sequence have developed a gaseous (and stellar) intra-group medium through physical interactions. The top sequence includes the possible precursors of dry mergers, as galaxies within these groups evolve separately and are likely to have a very low, or no gas content when they eventually merge. }\\label{fig:evo} \\end{figure}  The work presented here relates qualitative age dating of thousands of star clusters in seven Hickson compact groups (HCGs) to the CG evolutionary diagram proposed by \\citet[Fig.~\\ref{fig:evo}]{isk10}. This draws upon results presented in an ongoing series of papers on individual groups \\citep{gallagher10,isk10,fedotov11}, and on the collective properties of HCGs \\citep{johnson07,gallagher08,walker10,tzanavaris10}. In brief, the diagram classifies galaxy groups according to a gas richness criterion. With the basic assumption that H\\,{\\sc i} represents the reservoir of gas available for further star formation (star formation potential), \\citet{johnson07} uses the ratio of gas mass to dynamical mass (total mass) to split CGs into three types: I, II, and III, for gas-rich, intermediate and gas-poor. This is then filtered through the spatial distribution of the H\\,{\\sc i} gas \\citep[cf.]{vm01}, to give rise to a two-pronged diagram, shown in Figure~\\ref{fig:evo}. One sequence maps CGs where the gas in contained wholly within the galaxies; the other includes groups that have started to build an intra-group medium through the release of gas during interactions. ",
        "conclusions": "\\label{sec:summary} We have used the optical colours, and therefore ages, of thousands of \\textit{HST}-selected star clusters to infer the properties of their host galaxies, members of compact groups. We have linked the distribution of cluster colours to the evolutionary state of their hosts and provide evidence in favour of the evolutionary sequence for CGs proposed by \\citet{isk10}. While the lack of $U$-band coverage prohibits precision age-dating, the dust deficiency of the CG environment facilitates the qualitative age-dating of clusters with $BVI$ only. The full study, including modelling of the cluster populations with inferences on star cluster evolution, will be presented in a future work.  \\small  %"
    },
    "1107/1107.1316_arXiv.txt": {
        "abstract": "We have conducted an LTE abundance analysis for SAO\\,40039, a warm post-AGB star whose spectrum is known to show surprisingly strong He\\,{\\sc i} lines for its effective temperature and has been suspected of being H-deficient and He-rich. High-resolution optical spectra are analyzed using a family of model atmospheres with different He/H ratios. Atmospheric parameters are estimated from the ionization equilibrium set by neutral and singly ionized species of Fe and Mg, the excitation of  Fe\\,{\\sc i} and Fe\\,{\\sc ii} lines, and the wings of the Paschen lines. On the assumption that the He\\,{\\sc i} lines are of photospheric and not chromospheric origin, a He/H ratio of approximately unity is found by imposing the condition that the adopted He/H ratio of the model atmosphere must equal the ratio derived from the observed He\\,{\\sc i} triplet lines at 5876, 4471 and 4713\\AA, and singlet lines at 4922 and 5015\\AA. Using the model with the best-fitting atmospheric parameters for this He/H ratio, SAO\\,40039 is confirmed to exhibit mild dust-gas depletion, i.e., the star has an atmosphere deficient in elements of high condensation temperature. The star appears to be moderately metal-deficient with [Fe/H]=$-$0.4 dex. But the star's intrinsic metallicty as estimated from Na, S and Zn, elements of a low condensation temperature, is $\\rm [Fe/H]_o$\\footnote{$\\rm [Fe/H]_o$ refers to the star's intrinsic metallicity}=$\\simeq -0.2$. The star is enriched in N and perhaps O too, changes reflecting the star's AGB past and the event that led to He enrichment. ",
        "introduction": "SAO\\,40039, also known as IRAS 05040+4820 and BD $+48^\\circ 1220$, is judged by several criteria to be a post-AGB star: its spectral type (A4Ia), its position in the IRAS color-color diagram and the double-peaked spectral energy distribution indicating a detached youthful cold dust shell around the central star. Using  VBRIJHK photometry and adopting a simple dust shell model, \\citet{Fujii2002} find a dynamical age of 2460 years for the dust shell with dust at 97K, and estimate the star to have a core mass of 0.55M$_{\\odot}$ indicating  a low-mass progenitor for this PAGB star.  \\citet{Kloch2007} have presented an  abundance analysis  of SAO\\,40039 based on  high-resolution spectra with a wavelength coverage of 4500-6760\\AA. They have also discussed the star's radial velocity changes and also time-dependent differential velocity shifts between different classes of absorption lines. Their spectra showed a variable H$\\alpha$ line with two emission components. Variable emission was also reported in H$\\beta$ and in some lines of Si\\,{\\sc ii}, Fe\\,{\\sc i} and Fe\\,{\\sc ii}. But, in particular, these authors pointed out the abnormal strength of the He\\,{\\sc i} 5876\\AA\\ absorption feature in this star with the `low' effective temperature of about 8000 K. Their abundance analyses conducted with model atmospheres computed for a normal He/H ratio (=0.1) gave the He/H ratio of 0.7.  Detection of He\\,{\\sc i} lines in absorption in several A and F-type PAGB stars has been reported previously. The 4471\\AA\\ line was found by \\citet{Waelk92} in HD 44179, the central star of the Red Rectangle. HD 187885, a PAGB star with  a high C/O ratio and an enrichment of $s$-process elements, shows the He\\,{\\sc i} 5876 \\AA\\ line \\citep{Van96}. \\citet{Klo2002} report five He\\,{\\sc i} lines in HD 331319. Detection of He\\,{\\sc i} lines in A and F-type supergiants suggest either a He enrichment, if the lines are of photospheric origin, or a contribution from a thick hot chromosphere of presumably a normal He abundance. The above studies adopt the assumption of a photospheric origin and report the photosphere to be moderately H-poor and He-rich. However, these analyses use model atmospheres computed for a normal He/H ratio and do not iterate on the construction of the model atmosphere and the  abundance analysis to obtain agreement between the input and output He/H ratio.  Our goal in making a fresh analysis of SAO 40039 is to iterate to find the self-consistent He/H ratio. In undertaking the analysis, we  also adopt the assumption that the He\\,{\\sc i} lines originate in a He-rich photosphere and are not chromospheric in origin. ",
        "conclusions": "Though, the existence of extremely H-deficient stars with He/H ratio of about 10,000, showing He\\,{\\sc i} lines in their absorption spectra, is known \\citep{Pan2001,Pand2006}, SAO\\,40039, from our estimated He/H ratio of about unity, appears to be the first detection of a mildly H-deficient star. \\citet{Lam96} had commented on the possibilities of such stars and on the difficulty in detection for stars too cool to show He\\,{\\sc i} lines. With an effective temperature of about 8000K, SAO\\,40039 seems to be responding to this call.  It is very likely that other post-AGB stars may belong to the mildly H-deficient class. Several examples like SAO 40039 have detectable He\\,{\\sc i} lines. The question then arises; are post-AGB stars of a temperature too low to provide He\\,{\\sc i} lines also He-rich and how can their H-deficiency be unmasked? Exploration of post-AGB stars showing He\\,{\\sc i} lines need to be continued with models of appropriate H/He ratios. Additionally, non-LTE abundance analyses as described by \\citet{Pandey2011} should be performed to further refine knowledge of the chemical compositions.  We are thankful to the VBO observing support team, in particular to Mr G. Selvakumar, for obtaining a spectrum of SAO 40039 using the VBT echelle spectrometer. DLL thanks the Robert A. Welch Foundation of Houston, Texas for support through grant F-634.   \\clearpage"
    },
    "1107/1107.4533_arXiv.txt": {
        "abstract": "{ The brilliant Balmer H$\\alpha$ line in the stationary emission spectrum of the Galactic microquasar SS 433 has a broad component ($\\sim$ 1000 km s$^{-1}$). This is formed in the wind blowing from the accretion disk of the compact object, which orbits the centre of mass of the binary at $\\sim$ 175 km s$^{-1}$ with a 13-day period. The centroid of the H$\\alpha$ emission line formed in the wind has an imperfect memory of the motion of its source.} {The aim is to understand how this emission line is left with a blurred memory of the motion of the source of the wind.} {We analysed stationary H$\\alpha$ spectra, taken almost nightly over two orbital periods of the binary system.} { The loss of memory by emission lines from the wind is easily understood quantitatively, if any parcel of wind that has been newly conditioned to radiate in H$\\alpha$ continues its emission for a period of days. If the decay is exponential, a mean lifetime of about two days is all that is required. This timescale is very close to what is required for the decay of H$\\alpha$ from the putative circumbinary disk to be responsible for the narrow components of the H$\\alpha$ emission line. It is also close to the timescales observed for the individual bolides that radiate H$\\alpha$ in the relativistic jets.} { The structure of the stationary H$\\alpha$ emission line is now well understood. The broad component is undoubtedly formed in the wind from the accretion disk and the two narrow components carry no such signature. Their properties are entirely consistent with an origin in a circumbinary disk. The characteristic lifetime of H$\\alpha$ emitting material is a few days, as observed for the relativistic bolides in the jets.} ",
        "introduction": "We consider a very luminous star blowing a dense wind and a particular shell within that wind. It is accelerated rapidly to the terminal velocity, perhaps several thousand km s$^{-1}$. When close to the terminal velocity that expanding shell is effectively detached from the star and thereafter continues to expand with its centre following a straight-line trajectory, the velocity vector of which is equal to the velocity vector of the source at the time of detachment. If the very luminous source is a member of a binary system, the trajectory of the centre of a shell of wind does not deviate as the star continues in its closed orbit. Thus if emission lines are formed in the shell after the time of detachment, the mean Doppler shift of the shell spectrum will reflect the velocity vector of the source at the time of detachment rather than at the time of formation of the emission lines. If any given shell glows for a time lasting a significant fraction of the binary orbit, at any one moment the centre of any emission line will reflect the velocity of the source, but delayed and averaged over part of the orbit. The centre of an emission line formed in the wind will (for an approximately circular orbit seen approximately edge on) execute a sinusoidal oscillation with a velocity amplitude smaller than that of a line formed in the atmosphere co-moving with the star. This scenario is speculative, but might be relevant to binary systems containing Wolf-Rayet stars or objects with at least some of their observational characteristics. One such object is the Galactic microquasar SS 433, and in this note I suggest that the emission spectrum associated with the wind from the accretion disk exhibits just this behaviour.   SS 433 is very luminous and famous for its continual ejection of plasma in two opposite jets at approximately one quarter the speed of light. The system is a 13-day binary ( probably powered by supercritcal accretion by the compact member from its companion), and the orbital speed of the compact object is fairly well established as $\\sim$ 175 km s$^{-1}$. Despite the fame of the emission lines from the relativistic jets, the so-called stationary emission lines are much more intense, in particular the brilliant H$\\alpha$ line. Most aspects of the system are reviewed in Fabrika (2004). The stationary lines in fact are not quite stationary, rather  they exhibit sinusoidal Doppler shifts with a 13-day period; the oscillation of H$\\beta$ and He I stationary emission lines first demonstrated that SS 433 is a binary (Crampton, Cowley \\& Hutchings 1980). The velocity amplitudes of the Balmer series and He I lines, when treated as though they come from a single source, are very roughly 70 km s$^{-1}$ and do not provide a direct measure of the orbital speeds of either the compact object or its companion, because they are redshifted most when the compact object is very close to eclipse by the companion (orbital phase 0), rather than a quarter of a period earlier when the compact object is receding fastest from us (orbital phase 0.75); see for example Crampton \\& Hutchings (1981) and Gies et al (2002). This phasing of the emission lines relative to the photometric ephemeris naturally suggests that the Balmer and He I emission lines are formed in an accretion stream passing from the donor to the companion. However, Gies et al (2002), evidently influenced by the extreme brilliance of the stationary H$\\alpha$, proposes that these emission lines are formed in the wind from the accretion disk, but at distances comparable to or greater than the dimensions of the binary system. In that paper it is suggested that interference with the wind from the disk by the companion, or indeed by a wind from the companion, creates a void in the disk wind. This void would enfeeble the blue part of the spectrum when the companion is between us and the compact object (and conversely enfeeble the red half a period later). The origins of the stationary emission lines were later clarified by analysis of the stationary H$\\alpha$ line, in terms of a superposition of several Gaussian profiles (Blundell, Bowler \\& Schmidtobreick 2008), exploiting a remarkable sequence of spectra taken nightly over several orbits (see Schmidtobreick \\& Blundell 2006).   \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=15cm]{wigglines1.eps} \\caption{ Data from Fig.1 of Blundell, Bowler \\& Schmidtobreick (2008) for the wavelengths of the centroids of three Gaussians fitted to the H$\\alpha$ spectra. The blueshifted narrow Gaussian is denoted by +, the redshifted by x, and the broad Gaussian (the H$\\alpha$ profile of the wind) by $*$. The centre of the wind profile wanders back and forth between the railroad tracks, which have been attributed to a circumbinary disk. Orbital phase 0 (primary eclipse) occurs at JD +255. } \\label{fig:railroad} \\end{center} \\end{figure} ",
        "conclusions": "The broad component of the stationary H$\\alpha$ line was identified as formed in a wind from the accretion disk of SS 433 on the basis of the line of sight speed varying with the nodding of the disk (Blundell, Bowler \\& Schmidtobreick 2008). The amplitude and phase of the motion of the centroid of this component have now been shown to be entirely consistent with this origin, provided the emission decay time within the wind is a few days. The broad component of the H$\\alpha$ is indeed formed in the wind on a scale comparable to or greater than that of the binary system (Gies et al 2002). In contrast, the two narrow components show no sign of an origin within the wind; indeed, the narrow components of both H$\\alpha$ and He I are not consistent with such an origin but are well described in terms of emission from a fairly close circumbinary disk, stimulated by the intense radiation from the vicinity of the compact object (Bowler 2010, 2011b). It is interesting that in that circumbinary disk model the decay time for H$\\alpha$ emission  must again be of the order of a few days. Thus there are two phenomena, not directly related, that are described very well by the phenomenological assumption that H$\\alpha$ emission from dense plasma has, following ignition, a decay time of the order of days. There remains the question of by what mechanisms H$\\alpha$ radiation is sustained for several days after ignition. That at least one mechanism exists enabling clumps of plasma to exhibit H$\\alpha$ decay times of a few days is certain: the individual bolides ejected at about one quarter the speed of light in the jets for which SS 433 is famous are visible in H$\\alpha$ for several days ( see for example Vermeulen et al 1993, Gies et al 2002, Blundell, Bowler \\& Schmidtobreick 2007 )."
    },
    "1107/1107.1905_arXiv.txt": {
        "abstract": "A statistical analysis of radio flare events was performed by using the event list of Nobeyama Radioheliograph in 1996-2009. We examined center-to-limb variations of 17GHz and 34GHz flux by dividing the flare events into different groups with respect to the `thermal plasma richness' (ratio of the peak flux of soft X-ray to non-thermal radio emissions) and the duration of radio bursts. It is found that peak flux of 17 and 34GHz tend to be higher toward the limb for thermal-rich flares with short durations. We propose that the thermal-rich flares, which are supposed to be associated with an efficient precipitation of high energy particles into the chromosphere, have a pitch angle distribution of non-thermal electrons with a higher population along the flare loop. ",
        "introduction": "Non-thermal radio emissions observed in the impulsive phase of solar flares are produced by the gyrosynchrotron mechanism, which depends on a number of physical parameters such as electron energy spectra, their pitch angle distribution, magnetic field strength, angle between line of sight and  the magnetic field (viewing angle), and the number of electrons. Therefore, it is difficult to determine those physical parameters uniquely only from the observed quantities of individual radio burst. Statistical analyses of radio bursts by using a number of flare events provide us a way to find mutual relationships between different quantities, and thus are useful to restrict the possible domain of those physical quantities of the radio source. The pitch angle distribution of accelerated electrons is of a crucial importance for the problem of particle acceleration in solar flares. A clue to know the pitch angle distribution of accelerated particles could be obtained from the center-to-limb variations of observed radio emissions, since relativistic electrons trapped in flare loops emit the radio waves to the direction of their velocity, and the viewing angle effect, i.e., center-to-limb variation of the flare emission, can be related to  the pitch angle distribution of accelerated electrons.  A number of authors have performed statistical studies of flare parameters and investigated their center-to-limb variations, e.g., Dodson et al. (1954), Akabane(1956), Kundu (1959), Hervey (1964), Scalise (1970), Matsuura \\& Nave (1970), Croom \\& Powell (1971), Castelli \\& Barron (1977), Kosugi (1985), Pogodin et al. (1996), and Silva \\& Valente (2002). Some authors classified radio bursts based on the properties of their light curves, i.e.,  single/multiple bursts, gradual rise and fall type burst and so on. Dodson et al. (1954) showed that gradual rise and fall bursts in 2.8GHz has a high concentration toward the central part of the solar disk. In frequencies lower than 10GHz where the gyrosynchrotron emission is optically thick, a weak center-to-limb variation of the peak intensity was observed, i.e., the peak intensity gets slightly larger toward the limb. Takakura \\& Scalise (1970) and Ramaty (1969) argued that such center-to-limb variation can be explained with a bipolar magnetic field configuration containing non-thermal electrons with anisotropic pitch angle distribution, while the Razin Suppression in ambient plasmas taken into account to suppress the degree of the center-to-limb variation.   In 19 and 17GHz where the plasma is mostly optically thin, Croom \\& Powel (1971) and Kosugi (1985) found no significant center-to-limb variation of the peak flux. Contrary, Silva \\& Valente (2002) found a significant center-to-limb variation of gyrosynchrotron emission with a higher intensity toward the limb in 3.1, 5.2, 8.4, 11.8, 19.6, 35.0 and 50GHz, and that the variation gets more prominent in higher frequency ranges, i.e., in optically thin emission produced by higher energy electrons, than in lower frequency ranges. Vestrand et al. (1987) reported that hard X-ray emission with energy higher than 25keV has a center-to-limb variation by which the intensity is larger and the power law index is harder near the limb than these near the disk center.  Datlowe et al. (1977), however, found that the emission is isotropic and there are no significant center-to-limb variation in X-ray intensity  below 20keV.   Since the angle between line of sight and the magnetic field near the footpoints of flare loops gets larger toward the limb, flare events near-the-limb are to be brighter than the disk center events, if the footpoints are the major sources of the radio emission. In the theoretical view of Takakura \\& Scalise (1970), a center-to-limb variation of radio intensity is actually expected with a magnetic dipole field trapping non-thermal electrons, while its degree is dependent on the pitch angle distribution of the electrons. The fact that there found no significant center-to-limb variation in the frequency range below 17GHz was attributed  to the complexity and variety of the magnetic field configuration among flares (Kosugi 1985).  The aim of this paper is to find the insights on the pitch angle distribution of non-thermal electrons by investigating center-to-limb variations of radio intensity for different types of flares. To this end, we performed a statistical analysis by categorizing the flares in their duration, i.e., the impulsiveness. We use the radio data at two frequencies in optically thin ranges, 17 and 34GHz, to confirm the frequency dependence of the center-to-limb variation (Silva \\& Valente 2002). We also focus on the soft X-ray emission associated with the radio bursts. During the impulsive phase of solar flares, non-thermal electrons precipitate into the chromosphere, which leads to explosive plasma heating and pressure enhancement there. As a result, hot thermal plasmas emitting soft X-rays rapidly fill flare loops (these are so-called ``chromospheric evaporation\"). Since non-thermal electrons should play a significant role in generation of the thermal plasma as inferred from the empirical rule, so-called the Neupert Effect (Neupert 1968), the soft X-ray emission could be an indicator of the energy injection by non-thermal particles into the chromosphere. Furthermore, since the precipitation rate of non-thermal electrons would be related with the pitch angle distribution, the thermal plasma richness, i.e., the soft X-ray flux relative to the non-thermal flux, will provide a measure of the pitch angle distribution of electrons.  We use the radio burst data obtained by the Nobeyama Radioheliograph (NoRH) (Nakajima et al. 1994), and study the center-to-limb variation of intensity and other flare parameters. GOES data are also used to categorize the flares in their thermal richness. In Section 2, we describe the dataset of flare events used in this study. In Section 3, we define parameters for use of the statistical analysis, and describe their statistical properties. The `thermal emission index'  is introduced and relation to the other flare parameters are shown. In Section 4, we calculate correlation coefficients between each parameter and the heliocentric angle, and in Section 5 we discuss the results of the Section 4. ",
        "conclusions": "As shown, in Table \\ref{tab:17-hcd}, there is no center-to-limb variation in 17GHz peak flux when we take all events as a whole. This result agrees with Kosugi (1985). However, the peak flux of non-thermal radio bursts for flares with large TEI and short duration has a significant center-to-limb variation in both 17 and 34GHz ,i.e., are higher for near-the-limb events.  It is found in Table \\ref{tab:17-hcd}, \\ref{tab:34-hcd}, and Figure \\ref{fig:f5} (a)(b) that the higher frequency has stronger center-to-limb variation, as also can be seen in the power law index that gets harder toward the limb. This result agrees with the result of Silva \\& Valente (2002), i.e., the higher the frequency, the higher the tendency of the center-to-limb variation of flux.  The significant center-to-limb variation of radio flux found in thermal-rich events could be interpreted as follows; since thermal plasma is produced by the precipitation of non-thermal electrons to the chromosphere, it is supposed that high TEI flares have a pitch angle distribution with high population in small angle (along the magnetic field) so that an efficient precipitation of electrons to the chromosphere takes place. Under such pitch angle distribution, large number of electrons including lower energy ones can travel to the mirroring points located in a lower height in the corona, i.e. footpoints of the flare loop, where magnetic field is strong. Thus, lower energy electrons efficiently emit the gyrosynchrotron and the footpoints get the dominant sources of the flare emission. As the consequence, the observed radio emission gets brighter toward the limb.    The reason why only {\\it short} events have the center-to-limb variation of radio flux seems to be that, for {\\it long} events, soft X-ray flux is created by multiple non-thermal bursts, so that there is no one to one correspondence between the radio peak flux and the soft X-ray peak, thus TEI has no significant meaning and probably does not represent the pitch angle distribution of the non-thermal electrons. We also think that, in {\\it long} events, electrons tend to be trapped in coronal loop with a pitch angle distribution to be more perpendicular to the magnetic field, so that emission comes from looptop and there is less center-to-limb variation.  In conclusion, our statistical study of flare bursts presents the first observational evidence on that the efficiency of generating the thermal plasma is related to the pitch angle distribution of the non-thermal electrons of solar flares.  \\bigskip"
    },
    "1107/1107.1843_arXiv.txt": {
        "abstract": "We present an analysis of the atmospheric content of aerosols measured at Observatorio del Roque de los Muchachos (ORM; Canary Islands). Using a laser diode particle counter located at the Telescopio Nazionale Galileo (TNG) we have detected particles of 0.3, 0.5, 1.0, 3.0, 5.0 and 10.0 $\\mu$m size.\\\\ The seasonal behavior of the dust content in the atmosphere is calculated. The Spring has been found to be dustier than the Summer, but dusty conditions may also occur in Winter.\\\\ A method to estimate the contribution of the aerosols emissivity to the sky brightness in the near-infrared (NIR) is presented. The contribution of dust emission to the sky background in the NIR has been found to be negligible comparable to the airglow, with a maximum contribution of about 8-10\\% in the $K_{S}$ band in the dusty days. ",
        "introduction": "Superb observing conditions are crucial in order to obtain the best scientific output form the next generation of ground based telescopes, and this requires monitoring all relevant parameters that may affect observations. Due to their potentially detrimental impact on astronomical observations, atmospheric aerosols are among the most important parameters to be monitored in modern site testing campaigns. The performance and the safety of telescopes depend on the presence of atmospheric dust, which may deposit on  mirrors, increase atmospheric extinction, and emits in the infrared (IR) bands thus increasing the sky brightness.  Several studies about the atmospheric radiative effects of mineral dust exist. Gelado et al. \\cite{gelado03} analysed the mean aerosols content at Gran Canaria (Canary Islands) in periods 1997-1998 and 2002-2003. A mean grain size of 0.6-4.9 $\\mu$m and  a large annual variability in both density and size distribution have been found.  Cuevas\\&Baldasano \\cite{cuevas09} provided information on the incidence of the dust loaded African air masses at Observatorio del Teide in Tenerife (Canary Islands) and at Observatorio del Roque de Los Muchachos (ORM) in La Palma (Canary Islands). The analysis has been made using the high quality observations performed by the Iza\\~na atmospheric observatory. The in situ measurements for the 2002-2008 period show that the air is only partially affected by some dust loaded by African air mass intrusion. A significant PM10 concentration has been found only above the 80 percentile (26 $\\mu$g m$^{-3}$) at Iza\\~na.  \\begin{table*} \\begin{center} \\caption[]{Main characteristics of the dust monitors.} \\begin{tabular}{r | l | l} \\hline & ABACUS TM301 & LASAIRII 310B\\\\ \\hline input flow rate  & 0.1 Cubic Foot Minute            & 1.0 Cubic Foot Minute\\\\ size channels    & 0.3, 0.5, 1.0, 5.0 $\\mu$m        & 0.3, 0.5, 1.0, 3.0, 5.0, 10.0 $\\mu$m\\\\ light source     & laser diode ($\\lambda = 780$ nm) & laser diode ($\\lambda = 780$ nm)\\\\ sample rate      & 1 data per minute                & 1 data per 2 hours\\\\ output           & RS-232                           & Ethernet\\\\ running time     & August 2001 $-$ December 2006    & March 2007 $-$ today (January 2010 for this paper)\\\\ \\hline \\end{tabular} \\label{counters} \\end{center} \\end{table*}  The first analyses of the optical properties of dust and their impact on astronomy were done by Murdin \\cite{murdin85}, Stickland et al. \\cite{stickland87}, Guerrero et al. \\cite{guerrero98}, and Jimenez et al. \\cite{jimenez98}. More recently, Lombardi et al. \\cite{lombardi08} (hereafter Paper I)  analyzed 5+ years of time-series data spanning the period August 2001 to December 2006 obtained with the Telescopio Nazionale Galileo (TNG) dust counter, and showed that dust particles increase the extinction in the $B$, $V$ and $I$ bands at ORM.  In the present paper we extend the analysis of Paper I by adding almost 3 years of new data (March 2007 to January 2010) to cover a total period of almost 8 years. In the first part of the paper we calculate the seasonal trend of the atmospheric aerosol content on yearly and monthly basis. In the second part of the paper we present a method to estimate the thermal emissivity of the dust in the near-infrared (NIR). ",
        "conclusions": "Using a laser diode particle counter near the TNG at ORM  we have measured the densities of airborne aerosols of 0.3, 0.5, 1.0, 3.0, 5.0 and 10.0 $\\mu$m size. The seasonal trends of the particle content in the atmosphere have not changed significantly between 2001 and 2010.  The monthly distribution of aerosols is characterized by an increase during February-April and July-August of each year: the Spring is dustier than the Summer, but dusty conditions may also occur in Winter.  Using the Mie theory we calculated the dust absorption cross section and thus estimated the thermal emission in the 2MASS NIR spectral bands. Assuming that the dust particles are locally in thermal equilibrium with the surrounding air, and that the air temperature decreases with altitude with the wet vertical adiabatic lapse ($-0.006$ K m$^{-1}$), we found that the contribution of dust emission to the total sky background in the NIR is negligible comparable to the airglow component during both background and dusty conditions, with a maximum contribution of about 8-10\\% in the $K_{S}$ band in the dusty days."
    },
    "1107/1107.0768_arXiv.txt": {
        "abstract": "We present spatially resolved photometric and spectroscopic observations of two wide brown dwarf binaries uncovered by the SIMP near-infrared proper motion survey. The first pair  (SIMP\\,J1619275+031350AB) has a separation of $0.691\\arcsec$ (15.2 AU) and components T2.5+T4.0, at the cooler end of the ill-understood $J$-band brightening. The system is unusual in that the earlier-type primary is bluer in $J-K_{\\rm s}$ than the later-type secondary, whereas the reverse is expected for binaries in the late-L to T dwarf range. This remarkable color reversal can possibly be explained by very different cloud properties between the two components. The second pair (SIMP\\,J1501530-013506AB) consists of an L4.5+L5.5 (separation $0.96\\arcsec$, 30-47\\,AU) with a surprisingly large flux ratio ($\\Delta J =1.79$\\,mag) considering the similar spectral types of its components. The large flux ratio could be explained if the primary is itself an equal-luminosity binary, which would make it one of the first known triple brown dwarf systems. Adaptive optics observations could not confirm this hypothesis, but it remains a likely one, which may be verified by high-resolution near-infrared spectroscopy. These two systems add to the handful of known brown dwarf binaries amenable to resolved spectroscopy without the aid of adaptive optics and constitute prime targets to test brown dwarf atmosphere models. ",
        "introduction": "Binarity is relatively common among brown dwarfs (BDs) and recent work on these objects has shown that their properties differ from those of stellar binaries on many levels (see the review by \\citealt{Burgasser2007} and references therein). Their binarity fraction (10-30\\%) is lower than for Solar-type stars ($\\sim65\\%$, \\citealt{Duquennoy1991}) and follows the overall trend of decreasing binarity fraction with decreasing mass. Field brown dwarf binary systems are tighter than their more massive cousins, with few known systems with separations beyond 20\\,AU ($\\textless$2.3\\% of M7-L8 dwarfs have wide, 40 to 1000 AU, companions; \\citealt{Allen2007a}). Their mass ratios are also closer to unity (77\\% of systems have a mass ratio $\\ge 0.8$; \\citealt{Burgasser2007}). These different properties contain valuable information regarding their mechanism of formation. The dearth of wide BD binaries might be a signature of relatively violent formation processes: accreting proto-stellar cores that are ejected from the cluster before reaching a stellar mass \\citep{Reipurth2001} or photo-evaporation of pre-stellar cores by nearby massive stars \\citep{Whitworth2004}. Alternative processes have been proposed: gravitational collapse within a circumstellar disk \\citep{Boss2000} and collapse of very low mass molecular cloud cores \\citep{Boyd2005, Padoan2004, Padoan2005}. Observations of disks \\citep{Liu2003} and outflows around young BDs suggest that they undergo a T-Tauri-like phase, which points to a star-like formation mechanism \\citep{Jayawardhana2003, Mohanty2005, Scholz2006}. Each of these formation mechanisms leads to different BD binary frequency, separation distribution, and mass ratio distribution.  Beyond the interest in constraining BD formation models, binary BDs are natural benchmark objects to test atmosphere and evolutionary models. Their equal metallicity, coevality, and common distance put strong constraints when estimating these parameters from models by fitting both objects simultaneously \\citep{Liu2010}. Binaries that challenge evolutionary models have been found, most notably the 2MASS J05352184-0546085 system discovered by \\citet{Stassun2006} that shows a temperature reversal (the least massive object being the warmer one) that has tentatively been explained by widely different magnetic activity in the two components \\citep{Reiners2007}.  Binary BDs overlapping with the $J$-band brightening (roughly late-L to mid-T) are of particular interest as theoretical models cannot easily reproduce the observed evolution of BDs in this interval. The $J$-band brightening is partially explained by the presence of increased binarity at the L/T transition \\citep{Liu2006}. This is particularly the case for objects with the strongest brightening. However, even taking into account confirmed and suspected binaries, a smaller, but still significant, brightening is observed. This is also confirmed by the existence of binaries with flux reversals \\citep{Liu2006, Looper2008}. This brightening is probably best understood as resulting from the disappearance of dust-bearing clouds from the photosphere, but this temperature interval has proved challenging to describe with a self-consistent atmosphere model (\\citealt{Marley2010} and references therein). Binaries, especially those amenable to spatially resolved spectroscopy, provide unique constraints to these models. Tighter binaries have orbital periods short enough for orbital monitoring and provide further constraints on models through dynamical masses \\citep{Liu2008, Dupuy2009, Konopacky2010}.  In 2006, we have undertaken a near-infrared proper motion survey ({\\it Sondage Infrarouge de Mouvement Propre} - SIMP; \\citealt{Artigau2009}) with the wide-field near-infrared camera CPAPIR \\citep{Artigau2004} at the CTIO 1.5\\,m and OMM 1.6\\,m \\citep{Racine1978} telescopes and have covered $\\sim35\\%$ of the sky up to now. BD candidates were found using both the SIMP and 2MASS database obtained at different epochs to identify high proper motion sources. Spectroscopic follow-up of high proper motion candidates has been done with GNIRS \\citep{Elias2006}, SpeX \\citep{Rayner2003} and NIRI \\citep{Hodapp2003} in 2006, 2007 and 2008. More than 80 new L dwarfs and 14 new T dwarfs (Artigau et al., in prep.) were confirmed.  We did not expect to detect any resolved binaries in our sample given that the  binary BDs separation distribution peaks at $\\sim4$\\,AU \\citep{Maxted2005, Burgasser2007} and that all of our observations were seeing-limited. We nevertheless searched systematically through our spectroscopy acquisition images and found two partially resolved binaries: a pair of mid-L dwarfs separated by $\\sim1.0\\arcsec$ (SIMPJ1501530-013506, hereafter {\\simpbinl}) and a pair of mid-Ts separated by $\\sim0.7\\arcsec$ (SIMPJ1619275+031350, hereafter {\\simpbint}). Following these discoveries, we obtained a series of observations to characterize the individual components of these systems.  In \\textsection~\\ref{spectro} we present the NIRI resolved spectroscopy of both systems. In \\textsection~\\ref{niri_photometry}, \\textsection~\\ref{megacam_photometry}, \\textsection~\\ref{cpapir_photometry} and \\textsection~\\ref{lgs} we detail resolved and unresolved, seeing-limited, $i$ through $K_{\\rm s}$ photometric measurements and Laser Guide Star (LGS) AO observations of both binaries. Finally, the spectral typing of all components is discussed in \\textsection~\\ref{typing} and \\textsection~\\ref{discussion} details the properties of both systems. ",
        "conclusions": "We reported the discovery and characterization of two widely separated brown dwarf binaries: {\\simpbinl} and {\\simpbint}. Based on photometric measurements, {\\simpbinl} is most likely to be a triple system, with {\\simpbinl}A being a near-equal luminosity mid-L dwarf and {\\simpbinl}B a slightly later mid-L. This system has not been resolved as a triple in AO observations. This could be due to small physical projected separation at the time of observations and hope remains that future observations may resolve the suspected inner binary. It is also be possible that the semi-major axis is too small ($<$2\\,AU) to be resolved with current systems. In that case the system's binarity could be tested through very-high resolution spectroscopy to detect double lines or through still higher spatial resolution observation with either non-redundant mask imaging on 10-m class telescopes or future larger observatories. {\\simpbinl} is expected to have a $\\textgreater 475$\\,yr orbital period (total mass $\\textless0.12$M$_\\odot$). If {\\simpbinl}A proves to be an unresolved binary, then the expected orbital period for the outer component will be even larger; $\\textgreater 750$\\,yr for a total system mass of $\\textless0.18$M$_\\odot$.  The \\simpbint\\ components also display odd relative photometric properties: a near-infrared color reversal unlike anything seen in known L or T binaries. The toy model we presented shows that this inversion can be explained by the two components of the binary having very different dust content. This scenario may be tested with a more complete characterization of the SED through mid-infrared imaging and far-red spectroscopy. The relatively wide separation of the system implies that under good observing conditions (i.e. seeing better than about $0.5\\arcsec$), resolved spectroscopy can be obtained without the aid of adaptive optics, simplifying follow-up work. {\\simpbint} is expected to have a $\\textgreater 175$\\,yr period (total mass $\\textless0.12$M$_\\odot$). Neither system will provide significant constraints on dynamical masses within a useful time baseline.    \\begin{table}[!tbp] {\\tiny \\begin{center} \\caption{Observations log\\label{tbl-3}} \\begin{tabular}{lcccc} \\hline\\hline Instrument & Telescope & date          & filter & integration \\\\ \\hline \\multicolumn{5}{c}{\\simpbinl}\\\\ \\hline NIRI$^{\\rm a,b}$     & Gemini North& 2008/04/28 & $J^{\\rm g}$  & 100\\,s\\\\ NIRI$^{\\rm a,b}$     & Gemini North& 2008/04/28 & $H^{\\rm h}$  & 100\\,s\\\\ NIRI$^{\\rm a,b}$     & Gemini North& 2008/04/28 & $K^{\\rm i}$  & 100\\,s\\\\ NIRI$^{\\rm a,c}$     & Gemini North& 2008/05/19 & $J$  & 1200\\,s\\\\ NIRI$^{\\rm a,c}$     & Gemini North& 2008/05/19 & $H$  & 1200\\,s\\\\ NIRI$^{\\rm a,c}$     & Gemini North& 2008/05/19 & $K$  & 1200\\,s\\\\ NIRI LGS$^{\\rm a,d}$ & Gemini North& 2008/05/23 & $J$  & 810\\,s\\\\ NIRI LGS$^{\\rm a,d}$ & Gemini North& 2008/05/23 & $H$  & 810\\,s\\\\ NIRI LGS$^{\\rm a,d}$ & Gemini North& 2008/05/24 & $K_{\\rm s}$& 810\\,s\\\\ CPAPIR$^{\\rm e}$     &OMM& 2006/04/11 & $J$  &  24\\,s\\\\ % CPAPIR$^{\\rm e}$     &OMM& 2009/02/17 & $J$  &  1201\\,s\\\\ % CPAPIR$^{\\rm e}$     &OMM& 2009/02/17 & $H$  & 1169\\,s\\\\ % CPAPIR$^{\\rm e}$     &OMM& 2009/03/03 & $K_{\\rm s}$&   568\\,s\\\\ % MegaCam$^{\\rm e}$ &CFHT& 2006/07/20 & $i$ & 200\\,s\\\\ \\hline \\multicolumn{5}{c}{\\simpbint}\\\\ \\hline NIRI$^{\\rm a,b}$     & Gemini North& 2008/05/18 & $J^{\\rm g}$  &  900\\,s\\\\ NIRI$^{\\rm a,b}$     & Gemini North& 2008/05/18 & $H^{\\rm h}$   & 1170\\,s\\\\ NIRI$^{\\rm a,b}$     & Gemini North& 2008/05/18 & $K^{\\rm i}$   &  900\\,s\\\\ NIRI$^{\\rm a,c}$     & Gemini North& 2008/05/03 & $J$  &  100\\,s\\\\ NIRI$^{\\rm a,c}$     & Gemini North& 2008/05/03 & $H$  &  100\\,s\\\\ NIRI$^{\\rm a,c}$     & Gemini North& 2008/05/03 & {\\Ks}  &  100\\,s\\\\ NIRC2$^{\\rm d}$      &Keck-II& 2008/04/01 & $J$  &      540\\,s     \\\\ NIRC2$^{\\rm d}$      &Keck-II& 2008/04/01 & $H$  &    540\\,s      \\\\ NIRC2$^{\\rm d}$      &Keck-II& 2008/04/01 & {\\Ks}&    540\\,s       \\\\ NIRC2$^{\\rm d}$      &Keck-II& 2008/05/29 & CH$_4$s  &    200\\,s      \\\\ NIRC2$^{\\rm d}$      &Keck-II& 2008/05/29 & {\\Ks}&    360\\,s       \\\\ CPAPIR$^{\\rm e}$     &OMM& 2006/08/21 & $J$  &  24\\,s\\\\ % CPAPIR$^{\\rm e}$     &OMM& 2009/03/04 & $H$  &   601\\,s\\\\ % CPAPIR$^{\\rm e}$    &OMM & 2009/03/04 & $J$  &  1039\\,s\\\\  % CPAPIR$^{\\rm e}$     &OMM& 2009/03/09 & {\\Ks}&  1039\\,s\\\\  % MegaCam$^{\\rm f}$&CFHT & 2006/08/29 & $i$ &  200\\,s \\\\ MegaCam$^{\\rm f}$ &CFHT& 2009/06/01 & $z$ &  100\\,s \\\\ \\hline\\hline \\end{tabular} \\tablenotetext{a}{Gemini program ID : GN-2008A-Q-57.} \\tablenotetext{b}{Seeing limited, resolved spectroscopy.} \\tablenotetext{c}{Seeing limited, resolved photometry.} \\tablenotetext{d}{LGS-AO imaging.} \\tablenotetext{e}{Unresolved astrometry and photometry.} \\tablenotetext{f}{Resolved photometry.} \\tablenotetext{g}{NIRI ``$J$'' spectroscopy setup, useful domain : $0.985\\mu$m - $1.35\\mu$m.} \\tablenotetext{h}{NIRI ``$H$'' spectroscopy setup, useful domain : $1.36\\mu$m - $1.89\\mu$m.} \\tablenotetext{i}{NIRI ``$K$'' spectroscopy setup, useful domain : $1.90\\mu$m - $2.49\\mu$m.}   \\end{center} } \\end{table}   \\begin{table}[!htb] {\\tiny \\begin{center}   \\caption{Parameters of {\\simpbint} \\label{tbl-1}} \\begin{tabular}{lccc} \\hline\\hline Parameter& AB & A & B\\\\ \\hline 2MASS designation$^{\\rm a}$&J16192751+0313507 & & \\\\ $\\alpha^{\\rm a}$&$244.864664$&&\\\\ $\\delta^{\\rm a}$&$+03.230753$&&\\\\  $i[AB]_{\\rm MegaCam}$ & $22.50\\pm0.09$&$23.11\\pm0.14$&$23.41\\pm0.19$\\\\  $z[AB]_{\\rm MegaCam}$ & $19.17\\pm0.03$&$19.74\\pm0.05$&$20.15\\pm0.20$\\\\  $J^{\\rm a} _{\\rm 2MASS}$&$15.454\\pm0.071$ & &\\\\ $J_{\\rm 2MASS}^{\\rm b}$  &  $15.415\\pm0.009$ & $15.851\\pm0.013$ & $16.616\\pm0.013$\\\\ $J_{\\rm MKO}^{\\rm b}$    &  $15.207\\pm0.008$ & $15.652\\pm0.012$ & $16.390\\pm0.012$\\\\  $H^{\\rm a} _{\\rm 2MASS}$&$15.022\\pm0.098$&&\\\\ $H_{\\rm 2MASS}^{\\rm b}$  &  $14.916\\pm0.012$ & $15.482\\pm0.015$ & $15.894\\pm0.015$\\\\ $H_{\\rm MKO}$    &  $14.974\\pm0.012$ & $15.543\\pm0.015$ & $15.948\\pm0.015$\\\\  {\\Ks}$^{\\rm a}_{\\rm 2MASS}$&$15.015\\pm0.123$&&\\\\ $K_{\\rm 2MASS}^{\\rm b}$  &  $14.925\\pm0.012$ & $15.493\\pm0.015$ & $15.901\\pm0.015$\\\\ $K_{\\rm MKO}^{\\rm b}$    &  $14.868\\pm0.012$ & $15.432\\pm0.015$ & $15.848\\pm0.015$\\\\  $\\Delta J$ NIRC2&$0.738\\pm0.009$&&\\\\ $\\Delta J$ NIRI &$0.718\\pm0.008$&&\\\\  $\\Delta H$ NIRC2&$0.405\\pm0.008$&&\\\\ $\\Delta H$ NIRI &$0.388\\pm0.012$&&\\\\  $\\Delta {\\Ks}$ NIRC2&$0.312\\pm0.015$&&\\\\ $\\Delta {\\Ks}$ NIRI &$0.264\\pm0.005$&&\\\\  Separation&$0.691\\pm0.002\\arcsec$&&\\\\ Position angle&$71.23\\pm0.23\\,^\\circ$&&\\\\ $\\mu_\\alpha \\cos\\delta$&$70\\pm7$\\,mas/yr&&\\\\ $\\mu_\\delta$&$-289\\pm9$\\,mas/yr&&\\\\ H2O-J$^{\\rm c}$&&$0.510$ $[\\rm{T}2.4]$&$0.415$ $[\\rm{T}3.5]$\\\\ CH4-J$^{\\rm c}$&&$0.612$ $[\\rm{T}2.7]$&$0.542$ $[\\rm{T}3.6]$\\\\ H2O-H$^{\\rm c}$&&$0.502$ $[\\rm{T}2.4]$&$0.415$ $[\\rm{T}4.3]$\\\\ CH4-H$^{\\rm c}$&&$0.786$ $[\\rm{T}3.1]$&$0.590$ $[\\rm{T}4.1]$\\\\ CH4-K$^{\\rm c}$&&$0.532$ $[\\rm{T}3.1]$&$0.374$ $[\\rm{T}3.9]$\\\\ K/J $^{\\rm c}$&&$0.172$&$0.182$\\\\ Near-IR SpT&&$\\rm{T}2.5\\pm0.5$&$\\rm{T}4.0\\pm0.5$\\\\ Photometric distance$^{\\rm d}$&$22\\pm3$\\,pc&$22\\pm3$\\,pc&$22\\pm3$\\,pc\\\\ Physical separation$^{\\rm d}$&$15.4\\pm2.1$\\,AU&&\\\\ Temperature$^{\\rm e}$&&$1219\\pm100$\\,K&$1164\\pm100$\\,K\\\\ \\hline\\hline  \\end{tabular} \\tablenotetext{a}{From the 2MASS point source catalogue \\citep{Skrutskie2006}.} \\tablenotetext{b}{Combining the unresolved CPAPIR photometry with the NIRC2 flux ratios.} \\tablenotetext{c}{Index as defined by \\citet{Burgasser2006}.} \\tablenotetext{d}{Using the M$_{K_{\\rm s}}$ versus spectral type relation by \\citet{Liu2006}. Considering the peculiar colors of $J$-band flux ratio, we preferred M$_{\\Ks}$ to the M$_J$ polynomial relation for distance estimation. Uncertainty in the photometric distance include the {\\Ks}-band photometric uncertainty and an assumed 0.2\\,mag scatter in the M$_{\\Ks}$ versus spectral type relation.} \\tablenotetext{e}{Using \\citet{Stephens2009} polynomials; assumes spectral types of T2.7 and T3.9 (i.e. arithmetic means of spectral types derived from individual indices).}  \\end{center} } \\end{table}    \\begin{table}[!htb] {\\tiny \\begin{center} \\caption{Parameters of \\simpbinl\\label{tbl-2}} \\begin{tabular}{lccc} \\hline\\hline Parameter &AB& A & B\\\\ \\hline 2MASS designation$^{\\rm a}$&J15015302-0135068&&\\\\ SDSS designation&J150153.00-013507.2&&\\\\  $\\alpha$$^{\\rm a}$&225.470943&&\\\\ $\\delta$$^{\\rm a}$&-01.585236&&\\\\  $r_{\\rm SDSS}$&$22.706\\pm0.248$&&\\\\ $i_{\\rm SDSS}$&$20.802\\pm0.071$&&\\\\ $z_{\\rm SDSS}$&$18.622\\pm0.048$&&\\\\  $i[AB]_{\\rm MegaCam}$ & $20.44\\pm0.03$&&\\\\  $J^{\\rm a}_{\\rm 2MASS}$&$16.081\\pm0.094$&&\\\\ $J_{\\rm 2MASS}^{\\rm b}$  &  $15.978\\pm0.027$ & $16.168\\pm0.045$ & $17.960\\pm0.045$\\\\ $J_{\\rm MKO}^{\\rm b}$    &  $15.843\\pm0.027$ & $16.033\\pm0.045$ & $17.827\\pm0.045$\\\\  $H^{\\rm a} _{\\rm 2MASS}$&$14.952\\pm0.071$&&\\\\ $H_{\\rm 2MASS}^{\\rm b}$  &  $14.895\\pm0.015$ & $15.101\\pm0.019$ & $16.800\\pm0.019$\\\\ $H_{\\rm MKO}^{\\rm b}$    &  $14.956\\pm0.015$ & $15.163\\pm0.019$ & $16.858\\pm0.019$\\\\  {\\Ks}$^{\\rm a}_{\\rm 2MASS}$&$14.257\\pm0.086$&&\\\\ $K_{2MASS}^{\\rm b}$  &  $14.122\\pm0.028$ & $14.332\\pm0.034$ & $16.010\\pm0.034$\\\\ $K_{MKO}^{\\rm b}$    &  $14.182\\pm0.027$ & $14.393\\pm0.033$ & $16.063\\pm0.033$\\\\  $\\Delta J_{NIRI}$&$1.79\\pm0.04$&&\\\\ $\\Delta J_{NIRI LGS}$ &$1.77\\pm0.05$&&\\\\  $\\Delta H_{NIRI}$ &$1.70\\pm0.01$&&\\\\ $\\Delta H_{NIRI LGS}$&$1.67\\pm0.05$&&\\\\  $\\Delta K_{s NIRI}$&$1.67\\pm0.02$&&\\\\ $\\Delta K_{s NIRI LGS}$ &$1.62\\pm0.05$&&\\\\  Separation&$0.96\\pm0.01\\arcsec$&&\\\\ Position angle&$331.4\\pm0.1\\,^\\circ$&&\\\\  $\\mu_\\alpha \\cos\\delta$&$-225\\pm10$\\,mas\\,yr$^{-1}$&&\\\\ $\\mu_\\delta$&$-69\\pm12$\\,mas\\,yr$^{-1}$&&\\\\  H$_2$O$^A$  $^{\\rm c}$   &&$0.582$[L4.5]&$0.541$[L6.0]\\\\ H$_2$O$^B$  $^{\\rm c}$   &&$0.673$[L4.0]&$0.672$[L4.0]\\\\ H$_2$O$^C$  $^{\\rm c}$   &&$1.697$      &$1.788$\\\\ H$_2$O$^D$  $^{\\rm c}$   &&$0.673$      &$0.706$\\\\  H2O 1.5$\\mu$m$^{\\rm d}$   &&$1.520$[L4.5]&$1.482$[L3.5]\\\\ CH$_4$ 2.2$\\mu$m$^{\\rm d}$&&$0.987$[L5.0]&$1.114$[L8.0]\\\\  Near-IR SpT&&L$4.5\\pm0.5$&L$5.5\\pm1.0$\\\\ Photometric distance$^{\\rm e}$ && $31\\pm5$ / $44\\pm7$\\,pc$^{\\rm f}$  &  $56\\pm8$\\,pc\\\\ Physical separation&$30\\pm5$\\,AU / $47\\pm7$\\,AU$^{\\rm g}$&&\\\\ \\hline\\hline  \\end{tabular} \\tablenotetext{a}{From the 2MASS point source catalogue \\citep{Skrutskie2006}. } \\tablenotetext{b}{Combining the unresolved CPAPIR photometry with the NIRI flux ratios. } \\tablenotetext{c}{Index as defined by \\citet{Reid2001}.} \\tablenotetext{d}{Index as defined by \\citet{Geballe2002}.}  \\tablenotetext{e}{Using the \\citet{Liu2006} M$_{K_{\\rm s}}$ versus spectral type relations.} \\tablenotetext{f}{Respectively assuming that {\\simpbinl}A is itself a single object and an equal luminosity binary.} \\tablenotetext{g}{Assuming a $31\\pm5$ and $49\\pm7$\\,pc distance respectively.} \\end{center} } \\end{table}   \\clearpage"
    },
    "1107/1107.0597_arXiv.txt": {
        "abstract": "We present the results of our study of spectral energy distributions (SEDs) of a sample of ten low- to intermediate-synchrotron-peaked blazars.  We investigate some of the physical parameters most likely responsible for the observed short-term variations in blazars.  To do so, we focus on the study of changes in the SEDs of blazars corresponding to changes in their respective optical fluxes. We model the observed spectra of blazars from radio to optical frequencies using a synchrotron model that entails a log-parabolic distribution of electron energies. A significant correlation among the two fitted spectral parameters ($a$, $b$) of  log-parabolic curves and a negative trend among the peak frequency and spectral curvature parameter, $b$, emphasize that the SEDs of blazars are fitted well by log-parabolic curves. On considering each model parameter that could be responsible for changes in the observed SEDs of these blazars, we find that changes in the jet Doppler factors are most important. ",
        "introduction": "Studies of the two types of blazars, BL Lac objects and flat spectrum radio quasars (FSRQs), have shown that their Spectral Energy Distributions (SEDs) are characterized by a double peaked luminosity structure \\citep[e.g.][]{ghisellini1997}. The BL Lacs and FRSQs classes are defined according to the absence or presence of strong broad emission lines in their optical/UV spectrum, respectively. The  shapes of these bumps are characterized in the log($\\nu F_{\\nu}$) vs.\\ log $\\nu$ plot by a smooth spectrum extending through several frequency decades. The low-energy peak is well explained by synchrotron emission from relativistic electrons in a jet closely  aligned to the line of sight, with bulk Lorentz factor, $\\Gamma$ of order of 10--100 \\citep[e.g.,][and references therein]{maraschi1992, ghisellini1993, hovatta2009}.  Moreover, the position of the low-energy peak leads to a further classification of the blazars into three categories, depending on the peak frequency of their synchrotron bump, $\\nu_{peak}^S$: low-synchrotron-peaked (LSP) sources with $\\nu_{peak}^S$ $<$ 10$^{14}$ Hz; intermediate-synchrotron-peaked (ISP) the sources, those with 10$^{14}$ Hz $<$ $\\nu_{peak}^S$ $<$ 10$^{15}$ Hz; and high-synchrotron-peaked (HSP)  sources, with $\\nu_{peak}^S$ $>$ 10$^{15}$ Hz. This scheme is an extension of the classification introduced by \\citet{padovani1995} for BL Lacs \\citep[see][for details]{abdo2010a}.   The high energy (hard X-ray and $\\gamma$-ray component  of the SED is usually explained as arising from inverse Compton scattering  of the same electrons producing the synchrotron emission. These electrons interact either with the synchrotron photons themselves \\citep[synchrotron self-Compton, SSC; e.g.,][]{marscher1985} or with external photons originating in the local environment (external Compton, EC). In the latter case soft photons can be produced directly by the accretion disk \\citep[e.g.][]{dermer1993}, or indirectly, for instance those reprocessed by the broad line region (BLR) \\citep[e.g.][]{sikora1994}, or by the dust torus \\citep{blazejowski2000}. Alternative hadronic models, where the $\\gamma$-rays are produced by high-energy protons, either via proton synchrotron radiation or via secondary emission from photo-pion and photo-pair-production reactions, have also been proposed \\citep[see][and the references therein for a review]{bottcher2007}.  The  shape of these bumps is characterised in the Log($\\nu F_{\\nu}$) vs. Log $\\nu$ plot by a smooth spectrum extending through several frequency decades. Below and above the peaks the spectrum can be approximated with simple power law profiles with spectral indices above and below the peaks \\citep{ghisellini1996}, and such power law spectra are naturally produced if the the emitting electrons follow a power law distribution of energy. So, the broad band SED of blazars can be well approximated by a simple parabolic function with logarithms of its variables \\citep[e.g.,][and references therein]{massaro2004, massaro2006, tramacere2007}.  The log-parabolic function is one of the simplest ways to represent curved spectra and under simple approximations, can be obtained via a statistical electron acceleration mechanism where the acceleration probability decreases with the particle energy \\citep{massaro2006}.  The emission from blazars is known to be variable at all wavelengths. The flux variability is often accompanied by spectral changes. These SED changes are very likely associated with changes in the spectra of emitting electrons that arise from differences in the physical parameters of the jet. Hence modelling of blazar broadband spectra  is required to understand the extreme conditions within the emission region.  Not only is the broadband SED crucial to this understanding, but variability information is also needed to allow us to describe how high emission states arise and how they differ from the low states. This type of study is most important in discriminating between models.  Since it is reasonably to assume that only a few parameters of the model will change significantly between different states of the same source, such comparative modelling of broadband  spectra allows us to put rather tight constraints on those model parameters that are likely to change, at least under the assumption that all other parameters are held fixed for the different model fits \\citep[e.g.][]{mukherjee1999, petry2000, hartman2001}.   In this paper, we focus on the study of changes in the  SEDs of BL Lacs and FSRQs corresponding to observed changes in their respective optical fluxes. The SED variations are expected to be produced by changes in the spectra of the emitting electrons which in turn arise from variations in the physical parameters of the emission region. Therefore we needed to first construct models and then investigate how the SEDs observed in different states might be explained through  changes in these parameters. More explicitly, we modelled the observed spectra of blazars from radio to optical frequencies using a synchrotron model with the emitting electrons following a log-parabolic distribution of energy. This allows us to try to estimate the factors responsible for short term optical flares or short-term variability (STV) in blazars. In this paper, we will strictly concentrate on the low energy hump (synchrotron emission) portion of blazar SEDs. A complete investigation of the entire broad band SEDs of these objects, including X-ray and $\\gamma$-ray data, will be performed in the future.   This paper is structured as follows. Section 2 provides a brief description of the multi-frequency data we employed. In Section 3, we discuss the SED modelling and Section 4 provides our results.  Our discussion is given in Section 5 and we present our conclusions in Section 6. ",
        "conclusions": "We have carried out the radio to optical through mm, sub-mm and IR SED studies of a sample of ten blazars including five BL Lacs and five FSRQs, eight of which are LSPs and the other two ISPs. We modelled the SEDs of blazars using synchrotron spectra with  log-parabolic distributions. We found a significant negative correlation among $\\nu_{Peak}$ and the curvature term $r$, which implies that the acceleration efficiency of emitting electrons is inversely proportional to energy itself; however this correlation can also be explained by  the momentum diffusion term in the solution of Fokker-Planck equation. Also, a significant correlation between the two spectral parameters $a$ and $b$, implies that the log-parabolic curve is likely to be related to the statistical acceleration of emitting electrons.  Of course our modelling has significant limitations.  We are assuming that only one-zone dominates the emission at any given time, and this is unlikely to be an excellent approximation.  Although we strived for data obtained simultaneously, this was often impossible to obtain; therefore any variations within the ``high\" and ``low\" states have been averaged over time bins ranging from days to two months. This lack of simultaneity could have vitiated our results but only appears to have added scatter to all of the Figures 4--9.  We considered each likely factor that could be responsible for the changes in observed SEDs of blazars. If the electron energy density evolution governs the SED changes then one should expect a correlation between change in peak intensity vs change in peak frequency. Since we do not observe any such correlation we consider that probably the evolution of electron energy density cannot be responsible for the observed SED changes. Also for our entire sample of blazars, changes in $<$R$_{mag}$$>$ are not correlated with the respective changes in $B$, so the change in magnetic field strength is probably not responsible for the SED changes we saw.  However, for the BL Lacs alone the SED changes may be driven by changes in $B$. We find that the change in Doppler factor is significantly correlated with the change in peak intensity as well as with $<$R$_{mag}$$>$. So it is reasonable to suggest that the change in Doppler factor (either bulk velocity or viewing angle) is the primary driver that governs the SED changes in short term  variability (STV) in blazars."
    },
    "1107/1107.3182_arXiv.txt": {
        "abstract": "{Self-annihilating dark matter gravitationally captured by the Sun could yield observable neutrino signals at current and next generation neutrino detectors. By exploiting such signals, neutrino detectors can probe the spin-dependent scattering of weakly interacting massive particles (WIMPs) with nucleons in the Sun. We describe a method how to convert constraints on neutrino fluxes to a limit on the WIMP-nucleon scattering cross section. In this method all neutrino flavors can be treated in a very similar way. We study the sensitivity of neutrino telescopes for Solar WIMP signals using vertex contained events and find that this detection channel is of particular importance in the search for low mass WIMPs. We obtain highly competitive sensitivities with all neutrino flavor channels for a Megaton sized detector through the application of basic spectral selection criteria. Best results are obtained with the electron neutrino channel. We discuss associated uncertainties and provide a procedure how to treat them for analyses in a consistent way.}  \\notoc ",
        "introduction": "There is compelling observational evidence for the existence of dark matter; however, its nature remains unknown. A prominent candidate that naturally arises in supersymmetric~\\cite{Martin:1997ns} and large extra dimension models~\\cite{Appelquist:2000nn} is a weakly interacting massive particle (WIMP)~\\cite{Steigman:1984ac}. WIMPs are predicted to have masses ranging from a few GeV to several TeV~\\cite{Bertone:2004pz}. They could be detected through a variety of possible signals, indirectly through the byproducts of self-annihilations or decays, directly by scattering effects with nuclei, or production at colliders.   The nucleon coupling of a WIMP in the extreme nonrelativistic limit is characterized by a scalar (spin-independent) and an axial vector (spin-dependent) term. In general the spin-dependent (SD) couplings for WIMPs on protons and neutrons are expected to be of the same order of magnitude, and very similar for the spin-independent (SI) case~\\cite{PhysRevD.40.3132}. Throughout the paper we assume that the WIMP-proton and WIMP-neutron scattering cross sections are identical and do not consider isospin violating interactions, which have been discussed elsewhere (see for example ref.~\\cite{Feng:2011vu}).  The spin-dependent and spin-independent terms behave very differently if coherence across the nucleus is taken into account; the nucleon contributions interfere constructively to enhance the WIMP-nucleus elastic cross section. The WIMP-nucleon cross sections for spin-independent interactions are generally orders of magnitude smaller than those allowed for the spin-dependent case. Constructive interference across the nucleus leads to a quadratic enhancement with the nucleon number $(\\sim A^2)$ in the spin-independent couplings. By exploiting this dependence through the use of high $Z$ elements, direct detection experiments such as CDMS-II~\\cite{Ahmed:2009zw}, XENON100~\\cite{Aprile:2011hi}, and EDELWEISS-II~\\cite{Armengaud:2011cy} have been able to tightly constrain the spin-independent WIMP-nucleon scattering cross section.   Limits on the spin-independent cross section extend to $10^{-7}~{\\rm pb}$ ($10^{-43}~{\\rm cm}^2$) for WIMPs mass around 50~GeV. For the spin-dependent case constraints from direct detection experiments, such as SIMPLE~\\cite{Felizardo:2010mi}, COUPP~\\cite{Behnke:2010xt}, PICASSO~\\cite{Archambault:2009sm} are much weaker and range down to a minimum of about $2.8 \\times 10^{-2}$~pb for a WIMP mass of 45~GeV. Better sensitivities are achieved by indirect searches with neutrino telescopes for this case. Depending on the assumed annihilation channel, limits for the spin-dependent case from Super-Kamiokande~\\cite{Desai:2004pq} and IceCube~\\cite{Abbasi:2009uz} range from $10^{-2}$~pb to $10^{-4}$~pb. These tight constraints are obtained by looking for neutrino signals from self-annihilating dark matter gravitationally captured by the Sun. The comparison becomes possible, as the capture process is initiated by the same interaction that direct detection experiments exploit: a nuclear recoil in which a WIMP transfers a fraction of its kinetic energy to a nucleon. For the commonly assumed case of equilibrium between capture rate, $\\Gamma_{\\rm C}$, and annihilation rate, $\\Gamma_{\\rm A}$, neutrino fluxes, $\\phi_{\\nu}^{f}$, from the Sun are directly dependent on the WIMP-nucleon scattering cross section. Hence, the neutrino flux from the Sun or any constraint on it can directly be linked to the WIMP-nucleon scattering cross section.  While previous work has focused on the problem of converting the WIMP-nucleon scattering cross sections to muon neutrino generated muon flux, $\\phi_{\\mu}^{f}$~\\cite{Kamionkowski:1994dp,Wikstrom:2009kw}, we here focus on the neutrino flux instead. Muon fluxes inherently contain a dependence on the detector location as well as assumptions on neutrino interactions and muon propagation through some medium. Using the neutrino flux itself has the benefit that it allows to treat all neutrino flavors in a very similar way and the muon neutrino induced muon flux is just one special case of it. The neutrino flux is directly related to the study of vertex contained events at neutrino detectors. This class of events is in particular relevant for precision measurements of neutrino flavor ratios and the extraction of neutrino spectral information. We discuss results in light of a generic neutrino detector. Results are applicable for example to Super-Kamiokande, IceCube/DeepCore, KAMLAND, and future detectors proposed for long baseline neutrino oscillation experiments such as Hyper-K~\\cite{Nakamura:2003hk} and LENA~\\cite{Wurm:2011zn}. For a detector size (diameter) that is on the same order or larger compared to the average length of a muon track created by a muon neutrino with energy of interest, vertex contained events will be the dominant part of the expected signals in the detector.  This paper is structured as follows: In section~\\ref{sec:FluxConversion} we provide a description how to compare a neutrino flux from the Sun to the spin-dependent WIMP-nucleon scattering cross section, $\\sigma^{SD}$. Our result is based on the DarkSUSY package version 5.0.4~\\cite{Gondolo:2004sc}, which allows one to obtain conversion factors for given halo models and supersymmetric model parameters. Our study is focused on the WIMP mass range favored by supersymmetry (10~GeV to 10~TeV), however results are completely model-independent and applicable to any model that produces WIMPs in this mass range. We discuss the prospects of observing WIMP signals at neutrino detectors in section~\\ref{sec:Results}. We use a general approach that utilizes vertex contained events of all active neutrino flavors. Results are quoted for a data size of 200~kton$\\cdot$years and 5~Mton$\\cdot$years. We calculate differential neutrino fluxes, compared to detector backgrounds from atmospheric neutrinos and give event spectra. In section~\\ref{sec:Uncertainty} we describe how to incorporate and treat uncertainties related to the flux conversion. In particular we address uncertainties due to the dark matter distribution, which is underlying to all searches. We study the equilibrium condition and discuss procedures for scenarios when it is not satisfied. We summarize and draw our conclusions in section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have studied the comparison of neutrino fluxes from the Sun with the spin-dependent WIMP-nucleon cross section using DarkSUSY. The use of conversion factors for neutrino fluxes compared to muon fluxes hold some distinct advantages. Inherently, they do not depend on muon propagation effects and allow us to treat all neutrino flavors in the same way. We have discussed a procedure for the treatment of systematic uncertainties associated with the flux conversion and investigated the impact of these uncertainties. While, associated uncertainties on the annihilation rate can be significant they are not limiting to measurements.  Using vertex contained events and spectral event information we have estimated the sensitivity for neutrino telescope with a corresponding datasets of 200~kton$\\cdot$years and 5~Mton$\\cdot$years. For WIMP masses below 100~GeV the sensitivity to the spin-dependent WIMP-nucleon scattering cross section extends below $10^{-4}$~pb. Vertex contained events out perform muon rates once the detector size is on the same order as the corresponding track length of the muon neutrino induced muon. Even for higher WIMP masses, competitive results can be achieved using relatively large opening angles around the Sun that could be used to compensate for limited angular resolution. Assuming a comparable angular resolution is achieved the electron neutrino channel generally outperforms other flavor channels."
    },
    "1107/1107.3461_arXiv.txt": {
        "abstract": "{Barium stars are moderately rare, chemically peculiar objects, which are believed to be the result of the pollution of an otherwise normal star by material from an evolved companion on the asymptotic giant branch (AGB). } {We aim to derive carbon, nitrogen, oxygen, and fluorine abundances for the first time from the infrared spectra of the barium red giant star HD 123396 to quantitatively test AGB nucleosynthesis models for producing barium stars via mass accretion.} {High-resolution and high S/N infrared spectra were obtained using the Phoenix spectrograph mounted at the Gemini South telescope. The abundances were obtained through spectrum synthesis of individual atomic and molecular lines, using the MOOG stellar line analysis program, together with Kurucz's stellar atmosphere models. The analysis was classical, using 1D stellar models and spectral synthesis under the assumption of local thermodynamic equilibrium. } {We confirm that HD 123396 is a metal-deficient barium star ([Fe/H] = $-$1.05), with $A(\\mathrm{C}) = 7.88$, $A(\\mathrm{N}) = 6.65$, $A(\\mathrm{O}) = 7.93$, and $A(\\mathrm{Na}) = 5.28$ on a logarithmic scale where $A(\\mathrm{H}) = 12$, leading to [(C+N)/Fe] $\\approx$ 0.5. The $A(\\mathrm{CNO})$ group, as well as the $A(\\mathrm{Na})$ abundances, is in excellent agreement with those previously derived for this star using high-resolution optical data. We also found $A(\\mathrm{F}) = 4.16$, which implies [F/O] = 0.39, a value that is substantially higher than the F abundances measured in globular clusters of a similar metallicity, noting that there are no F measurements in field stars of comparable metallicity. } {The observed abundance pattern of the light elements (CNO, F, and Na) recovered here as well as the heavy elements ($s$-process) studied elsewhere suggest that the surface composition of HD 123396 is well fitted by the predicted abundance pattern of a 1.5{\\it M$_{\\odot}$} AGB model star with Z = 0.001. Thus, the AGB mass transfer hypothesis offers a quantitatively viable framework.} ",
        "introduction": "Barium stars are chemically peculiar objects --- dwarfs, subgiants, and giants --- that show an excess of carbon and s-process elements in their atmosphere \\citep[e.g.][]{tomkin89,barbuy92,allen06a,smilj07}. Identification of stars showing strong enhancements of neutron capture elements that can be produced by the $slow$ neutron capture process (the $s$-process), as well as prominent CH, CN, and C$_{\\rm 2}$ molecular lines, dates back to \\citet{bidelman51}. Since then, many other authors have sought to measure abundances of different species in order to better constrain the possible mechanisms by which barium stars can be formed.  The first detailed analysis of the barium star HD 46407 was by \\citet{burbidge57}, who found overabundances of barium and many other elements relative to solar or normal giants. Carbon, nitrogen, and oxygen were first analysed in barium stars by \\citet[][HR 774]{tomkin79} and \\citet[][$\\zeta$ Cap]{smith84}, who demonstrated that C were roughly solar and O slightly deficient by 0.1 dex. Nitrogen, on the other hand, was found to be enhanced by 0.3 dex with respect to the Sun in the barium stars analysed. More recently, \\citet{allen06a} and \\citet{allen07} acquired high-resolution spectra of giant and dwarf barium stars and obtained chemical abundances for several heavy elements(Sr, Y, Zr, Ru, Ba, La, Ce, Pr, Nd, Sm, Hf, Pb), that are produced by neutron-capture nucleosynthesis \\citep[the $s$ and $r$ processes, see e.g.,][]{sneden08}. In summary, detailed abundance analyses of the barium stars reveal that the heavy elements produced by the $s$-process are enhanced by factors of 2 -- 30 with respect to normal giants, whereas C is overabundant by a factor of 3 in the most extreme cases. While barium and CH stars likely had mass transferred from a former AGB companion, the latter group show strong C$_{2}$ molecular bands and higher C/O abundance ratios when compared to the former \\citep[e.g.,][]{luck91,mcclure97b}. The general view is that CH stars are more metal poor than barium stars, though there is no clear boundary.  The $slow$ neutron capture process occurs in the advanced thermally-pulsing asymptotic giant branch (TP-AGB) phase of low and intermediate-mass stars \\citep[e.g.,][]{busso99}. For this reason significant enrichment of $s$-process elements are not expected in stars in earlier evolutionary phases. The discovery that all observed barium stars belong to spectroscopic binary systems \\citep{bohm80,mcclure80,mcclure83,mcclure84,bohm85,jorissen88,mcclure90,udry98a,udry98b} with eccentricities significantly lower than a sample of spectroscopically normal G and K giants \\citep{mcclure90,jorissen98} is evidence that the C and large $s$-process overabundances are the result of accretion of enriched material from a former TP-AGB companion rather than the result of internal nucleosynthesis. The accretion of material is the direct result of mass transfer via Roche Lobe overflow or stellar winds, and the low eccentricities arise from tidal circularization \\citep{boffin88,han95,karakas00,izzard10}. In this context, CNO elements, along with $s$-process elements, are key tracers of the nucleosynthesis in the previous AGB progenitor, so they provide crucial constraints on the formation mechanism. According to the mass-transfer hypothesis, the former AGB stars would now be optically invisible white dwarfs, and several investigations with the IUE satellite have found the characteristic UV excesses typical of white dwarfs \\citep[e.g.][]{bohm00a,bohm00b}.  Another unique tracer of AGB nucleosynthesis is the element F. Outside of the solar system, F abundances were first determined in the pioneering work of \\citet{jorissen92} who analysed the correlation between F and C abundances in a sample of AGB stars. While it was clear from that study that F can be synthesized in AGB stars, and it is produced in the He-intershell along with C and $s$-process elements \\citep{jorissen92,lugaro04}, the nucleosynthetic origin of F in the Galaxy is still somewhat of a mystery. To date, F abundances have been determined for giant stars in the Galactic and extragalactic field \\citep[e.g.,][]{jorissen92,cunha05,pandey06,schuler07,pandey08,cunha08,utt08,abia10}, in Galactic and extragalactic globular clusters \\citep{cunha03,smith05,yong08c,lebzelter08}, and in post-AGB stars and planetary nebulae \\citep{werner05,zhang05,otsuka08}. For barium stars in particular, F abundances have been measured for only two stars \\citep{jorissen92} and thus much work is required on both the theoretical and observational sides to improve not only our knowledge of the nature of barium stars but also the astrophysical environment under which F is synthesized.  HD~123396 is of particular interest because it has an iron abundance of [Fe/H]\\footnote{We have adopted the usual spectroscopic notation, where {\\it A(X) = log[n(X)/n(H)] +12} and {\\it [X/Y] = log[n(X)/n(Y)]$_{*}$ $-$ log[n(X)/n(Y)]$_{\\odot}$}, where {\\it n}(X) and {\\it n}(Y) are the number particle density of \"X\" and \"Y\" respectively.} $\\approx -1$, and is the most metal-deficient barium star from the \\citet{allen06a} sample. The low metallicity of HD~123396 makes it unique among barium stars as it is right at the low end of the barium star metallicity distribution, and at a metallicity that encompasses the metal-rich tail of the halo, as well as the metal-poor tail of the disk \\citep[e.g.,][]{luck91,junq01,allen06a}.  The binary nature of HD 123396 has not been definitively established. There is no evidence, as yet, for radial velocity variations. \\citet{jorissen05} employed an additional check for binarity involving the comparison between the Hipparcos and Tycho-2 proper motions. However, this technique has been successfully employed for stars with parallaxes higher than 5 mas, while HD 123396 has a parallax of 1.78 $\\pm$ 0.73 mas \\citep{vl07}.  Here in this study we present the first abundance estimates of C, N, O, and F for HD~123396 using spectra taken at infrared wavelengths using the high-resolution, near-infrared Phoenix spectrograph at Gemini-South. HD~123396 exhibits overabundances of heavy elements at the three $s$-element peaks: the light $s$-peak, which includes the elements Sr, Y, and Zr, and denoted by the ratio [ls/Fe]; the heavy $s$-peak, which includes Ba, Ce, Nd, and Sm, and is denoted by the ratio [hs/Fe]; and the ratio of [Pb/Fe]. This makes HD~123396 an excellent tracer of chemically peculiar objects at low metallicities. The abundances of the CNO elements and F need to be considered together with the $s$-process elements, because in the context of the AGB mass transfer hypothesis, C and F are produced together with the heavy elements. ",
        "conclusions": "Employing high-resolution and high S/N data acquired with the Gemini Observatory Phoenix spectrograph, we have derived CNO, F, Na, and Fe abundances in the metal-deficient barium red giant star HD~123396. F along with C, N, and s-process elements are important for constraining the formation processes of chemically peculiar objects.  The detailed chemical abundance analysis performed here indicates that the nucleosynthetic history of HD~123396 is explained well in the context of the AGB mass transfer hypothesis. Furthermore, we derive, for the first time, [F/O] of a field star (despite its chemical peculiarity) in the regime spanned by the GC data ([Fe/H] $\\approx$ $-$1). Our results confirm that the GC and field stars have different values of [F/O] at a given $A(\\mathrm{O})$ abundance. However, gathering more data is vitally important in order not only to draw firm conclusions on the nucleosynthetic origin of F in the Galaxy, but also to investigate how chemically peculiar objects are formed and how the different stellar populations (field versus GCs) are chemically related."
    },
    "1107/1107.4594_arXiv.txt": {
        "abstract": "Recent observation of the tidally-excited stellar oscillations in the main-sequence binary KOI-54 by the {\\it KEPLER} satellite provides a unique opportunity for studying dynamical tides in eccentric binary systems.  We develop a general theory of tidal excitation of oscillation modes of rotating binary stars, and apply our theory to tidally excited gravity modes (g-modes) in KOI-54.  The strongest observed oscillations, which occur at 90 and 91 times the orbital frequency, are likely due to prograde $m=2$ modes (relative to the stellar spin axis) locked in resonance with the orbit.  The remaining flux oscillations with frequencies that are integer multiples of the orbital frequency are likely due to nearly resonant $m=0$ g-modes; such axisymmetric modes generate larger flux variations compared to the $m=2$ modes, assuming that the spin inclination angle of the star is comparable to the orbital inclination angle.  We examine the process of resonance mode locking under the combined effects of dynamical tides on the stellar spin and orbit and the intrinsic stellar spindown. We show that KOI-54 can naturally evolve into a state in which at least one $m=2$ mode is locked in resonance with the orbital frequency. Our analysis provides an explanation for the fact that only oscillations with frequencies less than 90-100 times the orbital frequency are observed. We have also found evidence from the published {\\it KEPLER} result that three-mode nonlinear coupling occurs in the KOI-54 system. We suggest that such nonlinear mode coupling may explain the observed oscillations that are not harmonics of the orbital frequency. ",
        "introduction": "Tides play an important role in binary star systems and in star-planet systems. While numerous studies of tidal effects have been based on the so-called equilibrium tide theory, which parametrizes tidal dissipation by an effective tidal lag angle/time or tidal quality factor (e.g., Darwin 1880; Goldreich \\& Soter 1966; Alexander 1973; Hut 1981), the underlying physics of tides in fluid stars and planets involves dynamical excitations of waves and oscillations by the tidal force (see Ogilvie \\& Lin 2007 and Zahn 2008 for recent reviews). Tides in highly eccentric systems are particularly rich in their dynamical behavior, since wave modes with a wide range of frequencies can be excited and participate in the tidal interaction. Various aspects of dynamical tides in eccentric binaries have been studied by Lai (1996,1997), Kumar \\& Quataert (1998), Witte \\& Savonije (1999,2001), Willems et al.~(2003), Ivanov \\& Papaloizou (2004) and Papaloizou \\& Ivanov (2010).  Recent observations of the binary star system HD 187091 (KOI-54) by the {\\it Kepler} satellite provide a unique opportunity for studying dynamical tides in eccentric binaries. KOI-54 consists of two A stars (mass $M_{1,2}=2.32,\\,2.38\\,M_\\odot$ and radius $R_{1,2}=2.19,\\,2.33\\,R_\\odot$) in an eccentric ($e=0.8342$) orbit with period $P=41.805$~days (Welsh et al.~2011). The binary is nearly face-on with orbital inclination $i_{\\rm orb}=5.52^\\circ$. In addition to periodic brightening events caused by tidal distortion and irradiation of the two stars during their close periastron passages, the power spectrum of the {\\it Kepler} light curve revealed 30 significant (with a signal-to-noise ratio $\\go 7$) stellar pulsation modes.  The observed mode periods range from 45~hours to 11 hours, corresponding to the mode frequency $f_\\alpha$ ranging from $22.42\\,f_\\orb$ to $91\\,f_\\orb$ (where $f_\\orb=P^{-1}$ is the orbital frequency). Most interestingly, twenty-one of these mode frequencies are integer multiples of $f_\\orb$ (with the ratio $f_\\alpha/f_\\orb$ differing from an integer by $0.01$ or less). The two dominant modes have frequencies that are exactly 90 and 91 times $f_\\orb$, with the corresponding flux amplitudes of 297.7~$\\mu$mag and 229.4~$\\mu$mag, respectively.  While dynamical tides in massive-star binaries have been studied before (e.g., Zahn 1977, Goldreich \\& Nicholson 1989 for circular binaries; Lai 1996,1997, Kumar \\& Quataert 1998 and Witte \\& Savonije 1999,2011 for eccentric binaries), KOI-54 represents the first example where tidally excited oscillations are directly observed and therefore serves as an explict demonstration of dynamical tides at work in the system. As discussed in Welsh et al.~(2011), the observed oscillation modes are puzzling: over 20 of the observed modes are nearly exact integer multiples of the orbital frequency, yet several others are not. It is not clear why the dominant modes are so prominent, e.g., why modes with frequencies 90, 91, 44, and 40 times $f_{\\rm orb}$ are clearly visible, and yet modes with frequencies greater than $91f_{\\rm orb}$ and those less than $20f_{\\rm orb}$ appear to be absent.  The goal of this paper is to explain some of the observational puzzles related to KOI-54 and to develop the general theoretical framework for studying tidally-excited oscillations in eccentric binary systems.  Our paper is organized as follows. In Section 2, we derive the general equations for calculating the energies of tidally excited oscillation modes in an eccentric binary. Our theory improves upon previous (and less rigorous) works, and provides a clear relationship between the resonant mode energy and non-resonant mode energy.  In Section 3 we study the properties of non-radial g-modes relevant to the stars in the KOI-54 system and calculate the non-resonant mode energies -- these serve as a benchmark for examining the effect of resonances.  In Section 4 we present our calculations of the flux variation due to tidally-forced oscillations. We show that the observed flux variation in KOI-54 can be largely explained when a high-frequency mode is locked into resonance (with the mode frequency equal to $90f_{\\rm orb}$). In Section 5 we study the possibility of resonance locking. We show that the combination of the secular tidal orbital/spin evolution and the intrinsic spindown of the star (e.g., due to stellar evolution) may naturally lead to resonance locking of a particular mode. Our analysis demonstrates that (in the KOI-54 system) a mode with frequency around $90f_{\\rm orb}$ can be resonantly locked, while modes with higher frequencies cannot.  In Section 6 we discuss the origin of the observed modes in KOI-54 with frequencies that are not an integer multiple of $f_{\\rm orb}$, including the evidence of nonlinear mode coupling. We conclude in Section 7 with a discussion of future prospects and remaining puzzles. ",
        "conclusions": "We have shown that many properties of flux oscillations detected in the KOI-54 binary system by {\\it KEPLER} can be understood using the theory of dynamical tidal excitations of g-modes developed in this paper. In particular, our analysis and calculation of the resonance mode locking process, which is driven by dynamical tides and intrinsic stellar spindown, provides a natural explanation for the fact that only those modes with frequencies ($\\sigma_\\alpha$) less than about 90-100 times the orbital frequency ($\\Omega$) are observed. Our result suggests that the KOI-54 system is currently in a resonance-locking state in which one of the stars has a rotation rate such that it possesses a $m=2$ mode with frequency $\\sigma_\\alpha=90\\Omega$ and the other star has a similar mode with $\\sigma_{\\alpha'}=91\\Omega$ --- these modes produce the largest flux variations detected in KOI-54. Our analysis shows that the system can evolve into and stay in such a resonance locking state for relatively long time periods, and it is reasonable to observe the system in such a state.  Other less prominent flux variations at lower frequencies can be explained by the $m=0$ tidally forced oscillations, many of which may be enhanced by the resonant effect. We have also found evidence of nonlinear three-mode coupling from the published {\\it KEPLER} data and suggested that the nonlinear effect may explain the flux variabilities at non-harmonic frequencies.  In our study, we have used approximate quasiadiabatic mode damping rates. Obviously, it would be useful to repeat our analysis using more realistic mode damping rates, calculated with the full non-adiabatic oscillation equations. Our calculations of the flux variations of tidally induced oscillations have also highlighted the importance of accurate treatment of non-adiabatic effects in the stellar photosphere. We have only considered g-modes (modified by stellar rotation) in this paper. Our general theory allows for other rotation-driven modes, such as inertial/Rossby modes. It would be interesting to study tidal excitations of these modes as well as the nonlinear 3-mode coupling effect in the future.  More detailed modelling of the observed oscillations may provide useful constraints on the parameters of the KOI-54 system, particularly the stellar rotation rates and spin inclinations. On a more general level, KOI-54 may serve as a ``laboratory'' for calibrating theory of dynamical tides, which has wide applications in stellar and planetary astrophysics (see references in Section 1).  While we believe that our current theory provides the basis for understanding many aspects of the KOI-54 observations, some puzzles remain.  In our current interpretation, the most prominent oscillations (at $90\\Omega$ and $91\\Omega$) occur in the two different stars, each having a $m=2$ mode resonantly locked with the orbit. While one can certainly appeal to coincidence, it is intriguing that the strongest observed oscillations are at the two consecutive harmonics ($N=90,\\,91$) of the orbital frequency. More importantly, as discussed in Section 5.3.3, when both stars are involved in resonance locking, the effective $N_c$ (above which resonance locking cannot happen) is reduced. For example, using our canonical stellar parameters ($\\kappa =0.05$), we find $N_c\\simeq 131$ [see equation (\\ref{nc2})], but the effective $N_c$ would be reduced to $92$ if the two stars have identical modes in resonance locking. In the case where the resonant energy dissipation rates in the two stars differ by a factor of more than a few, the effective $N_c$ of one star would remain close to 131, while the effective $N_c$ for the other star would likely be reduced to a value below 90. In this case, it would be unlikely to find the binary system in the resonance-locking state involving both stars with $N,N'$ close to 90, because it is most likely for modes with $N$ just less than $N_c$ to be locked in resonance (see Section 5.4.2.). Thus, we find it likely that the energy dissipation rates are nearly equal (to within a factor of 2) in each star in the KOI-54 system.  Finally, we may use our results to speculate on the evolutionary history (and future) of the KOI-54 system. The star with $M_2=2.38 M_\\odot$ and radius $R_2 =2.33 R_\\odot$ has an age about $3\\times 10^8$ years, according to our stellar model generated by the MESA code. During the resonance locking phase, the orbital eccentricity decreases on a relatively short time scale (of order $\\sim 10^8$ years). The current large eccentricity of KOI-54 then suggests that resonance locking has not operated for a large fraction of the system's history. When resonance locking is in effect, the orbital evolution time scale is set by primarily by the spindown time scale, $t_{sd}$, and the value of $\\delta_N$ for the resonant mode [see equation (\\ref{domres})]. In the future, the stars will continue to expand into red giants and the spindown time scale will decrease. This will cause the orbital evolution time scale to correspondingly decrease, leading to rapid orbital decay and circularization of the system.  While our paper was under review,  a paper by Burkart et al. (2011) appeared on arXiv.  Burkart et al. carried out non-adiabatic mode calculations and also attribute most of the oscillations to $m=0$ modes, although they do not reach a definite conclusion on the source of the $N=90,91$ oscillations. They did not consider resonance locking (instead they consider the qualitatively different phenomenon of \\textquotedblleft pseudosynchronous locks\"), and appeared to attribute all resonances to random chance. They also showed that 3-mode coupling (see Section 6) is possible."
    },
    "1107/1107.4241_arXiv.txt": {
        "abstract": "The Kelvin-Helmholtz instability is well known to be capable of converting well-ordered flows into more disordered, even turbulent, flows.  As such it could represent a path by which the energy in, for example, bowshocks from stellar jets could be converted into turbulent energy thereby driving molecular cloud turbulence.  We present the results of a suite of fully multifluid magnetohydrodynamic simulations of this instability using the HYDRA code.  We investigate the behaviour of the instability in a Hall dominated and an ambipolar diffusion dominated plasma as might be expected in certain regions of accretion disks and molecular clouds respectively.  We find that, while the linear growth rates of the instability are unaffected by multifluid effects, the non-linear behaviour is remarkably different with ambipolar diffusion removing large quantities of magnetic energy while the Hall effect, if strong enough, introduces a dynamo effect which leads to continuing strong growth of the magnetic field well into the non-linear regime and a lack of true saturation of the instability. ",
        "introduction": "The Kelvin-Helmholtz (KH) instability is an important instability in almost any system involving fluids: it can occur anywhere that has a velocity shear.  In astrophysical plasmas the KH instability can provide the means of producing turbulence in a medium or the mixing of material between two boundary layers.  The KH instability has been studied in a variety of astrophysical systems, from solar winds \\citep{amerstorfer07, bettarini06, hasegawa04} and pulsar winds \\citep{bucciantini06} to thermal flares \\citep{venter06}. Due to its ability to drive mixing and turbulence, the KH instability has been considered relevant in protoplanetary disks \\citep{johansen06, gomez05}, accretion disks and magnetospheres \\citep{li04}, and other jets and outflows \\citep{baty06}. Generally speaking, the assumptions of ideal magnetohydrodynamics (MHD) have been used in order to simplify the system of equations to be solved. These assumptions are, however, not always valid.  Weakly ionised plasmas, for example, contain a large fraction of neutral particles as well as a number of charged particle fluids with differing physical characteristics.  Interactions between the various species can introduce non-ideal effects. Ambipolar dissipation and the Hall effect are two non-ideal effects that can greatly influence the development of the KH instability in a system by altering the dynamics of the plasma and the evolution of the magnetic field.  Astrophysical examples of such weakly ionised systems include dense molecular clouds \\citep[e.g.][]{ciolek} and accretion disks around young stellar objects \\citep[e.g.][]{wardle99disk}. In these systems, the relevant length scales are such that non-ideal effects can play an important role \\citep{wardle04a, downes09, downes11}.   Many authors have investigated the role of the KH instability in both magnetised and unmagnetised astrophysical flows \\citep[e.g.][]{frank96, mala96, hardee97num, downes98, keppens99}.  Most of these studies have investigated the KH instability in the context of either hydrodynamics or ideal MHD.  We know that non-ideal effects are important in molecular clouds at length scales below about 0.2\\,pc \\citep[e.g.][]{ois06, downes09} and hence it is of interest to explore the KH instability in the context of either non-ideal MHD or, preferably, fully multifluid MHD.  In more recent years the emphasis of KH studies has been on including non-ideal effects.  \\citet{keppens99} studied both the linear growth and subsequent nonlinear saturation of the KH instability using resistive MHD numerical simulations. The inclusion of diffusion allowed for magnetic reconnection and non-ideal effects were observed through tearing instabilities and the formation of magnetic islands.  \\citet{palotti08} also carried out a series of simulations using resistive MHD. They found that, following its initial growth, the KH instability decays at a rate that decreases with decreasing plasma resistivity, at least within the range of of resistivities accessible to their simulations. They also found that magnetisation increases the efficiency of momentum transport, and that the transport increases with decreasing resistivity.  \\citet{birk02} examined the case of a partially ionised dusty plasma, using a multifluid approach in which collisions could be included or ignored.  They found that collisions between the neutral fluid and dust particles could lead to the stabilisation of KH modes of particular wavelengths. The unstable modes led to a significant local amplification of the magnetic field strength through the formation of vortices and current sheets. In the nonlinear regime they observed the magnetic flux being redistributed by magnetic reconnection. It was suggested that this could be applicable to dense molecular clouds and have important implications for the magnetic flux loss problem \\citep{umebayashi90}.  A comprehensive study was carried out by \\citet{wiechen06} which demonstrated the effect of dealing with the plasma using a multifluid scheme. This study focused on the effect of varying the properties of the dust grains. The results of the simulations led to the conclusions that more massive dust grains have a stabilising effect on the system while higher charged numbers have a destabilizing effect. It was found that there is no significant dependence on the charge polarity of the dust.  In a linear study of stellar outflows \\citet{watson04} described how the charged and neutral fluids are affected differently by the presence of a magnetic field. This study is carried out using parameters chosen to reflect those of molecular clouds, and so is particularly relevant to our own study. The principal result of this paper is that for much of the relevant parameter space, neutrals and ions are sufficiently decoupled that the neutrals are unstable while the ions are held in place by the magnetic field. Since the magnetic field is frozen to the ionised plasma, it is not tangled by the turbulence in the boundary layer. The authors predict that with well-resolved observations, there should be a detectably narrower line profile in ionised species tracing the stellar outflow compared with neutral species, since ionised species are not participating in the turbulent interface with the ambient interstellar medium. The paper also includes a study of the growth rate of the instability.  It is found that at short length scales, the growth rate is well approximated by the growth rate of the hydrodynamic system. At larger scales and for super-Alfv\\'{e}nic flows, the fastest growing mode is equal to that of the ideal MHD case.   \\citet{shad08} carried out an analytical study of the Kelvin-Helmholtz instability in dusty and partially ionised outflows. They investigated primarily the effect of the presence of dust particles by varying their mass, charge and charge polarity. It was found that as the charge of the grain increased, the growth timescales also increased, implying a stabilising effect on the system. The stability of the system was also examined for dependence on the mass of the dust particles. It was found that for stronger magnetic fields, this did not affect the stability of the system. However, for weaker magnetic fields, the larger dust particles had a stabilising effect on the growing modes. This was in agreement with previous laboratory experiments \\citep{luo01} and numerical simulations \\citep{wiechen06}. Finally, as the magnetic field strength increased, the growth timescale of the unstable modes at a particular perturbation wavelength decreased.  By examining the combinations of the wavelength of the perturbation used, and the resultant growth timescales of the instability, \\citet{shad08} concluded that the Kelvin-Helmholtz instability is a possible candidate for causing the formation of some of the physical structures observed in molecular outflows from young stars.  In this paper we perform numerical simulations of the complete evolution of the Kelvin-Helmholtz instability in a weakly ionised, multifluid plasma including both its linear development, saturation and its subsequent behaviour.  We include the physics of the Hall effect, ambipolar diffusion and parallel resistivity in these simulations and we analyse the role each of the former two effects on the development of the instability.  The aim of this work is to investigate the influence of multifluid effects on the growth and saturation of the KH instability.  In order to develop a full understanding of the roles of the various non-ideal effects, in particular the Hall effect and ambipolar diffusion, we run simulations with parameters chosen to simulate very high, medium and very low magnetic Reynolds number systems and with parameters chosen to ensure ambipolar-dominated flows and Hall-dominated flows.  We focus on gaining an understanding of the general characteristics of the KH instability in weakly ionised flows.  A detailed study of this instability in the specific context of molecular clouds, and which is of interest from the point of view of turbulence generation by stellar outflows, is the subject of a future work.   In section \\ref{sec:num-app} we outline the numerical and physical model employed, in section \\ref{sec:analysis} we discuss how we analyse the results of our simulations while in section \\ref{sec:results} we detail the results in both the linear and non-linear regimes, separating out the effects of ambipolar diffusion and the Hall effect in order to more fully understand the influence of each. ",
        "conclusions": "We have presented the results of a suite of fully multifluid MHD simulations of the KH instability in weakly ionised fluids such as, for example, molecular cloud material.  Through varying the collision coefficients between the various charged species and the neutrals we were able to investigate systems in which ambipolar diffusion dominates the multifluid effects and ones in which the Hall effect dominates.  We validated our KH simulations through comparison of an % ideal MHD simulation with previously published results and performed resolution studies for each of these cases to ensure that our conclusions are not unduly effected by our numerical resolution.  Our findings can be summarised as follows: \\begin{itemize} \\item The multifluid effects do not significantly influence the linear growth rate of the instability. \\item Ambipolar diffusion dramatically reduces the energy associated with the perturbed magnetic field.  This happens through diffusion % for moderate ambipolar resistivity, but through both diffusion and decoupling of the magnetic field from the bulk flow for higher resistivity. \\item The Hall effect, as expected, rotates the magnetic field out of the initial $xy$-plane. \\item In contrast to both ambipolar dominated and ideal MHD flows, the Hall effect causes the system to fail to settle to a quasi-steady state after saturation of the instability. \\item For moderate Hall resistivity the perturbed magnetic field contains higher energy than in the ideal MHD case, presumably due to a field topology which impedes (numerical) reconnection. \\item For high Hall resistivity strong dynamo action is seen as energy associated with the magnetic field grows without any apparent signs of saturation. \\end{itemize}"
    },
    "1107/1107.1378_arXiv.txt": {
        "abstract": "The cold dark matter theory predicts triaxial dark matter haloes. The radial distribution of halo ellipticity depends on baryonic processes and the nature of dark matter particles (collisionless or collisional). Here we show that we can use lensing flexion ratios to measure the halo ellipticity as a function of radius. We introduce a weight function and study the relationship between the first and second order statistics of flexion ratios, both of which can be used to reduce the bias in the estimate of ellipticity. We perform numerical tests for our method, and demonstrate that it can reduce the bias and determine the halo ellipticity as a function of radius. We also point out that the minimum mean flexion ratio can be used to trace the centres of galaxy clusters. ",
        "introduction": "Structures of cluster- and galactic-sized dark matter haloes predicted by N-body simulations of Cold Dark Matter (CDM) have several important features, e.g. a nearly universal radial density profile for haloes over a wide range of masses \\citep{1996ApJ...462..563N,1997ApJ...490..493N}, and triaxial shapes \\citep{2002ApJ...574..538J,2004IAUS..220..421S}. These features are related to the nature of dark matter particles, the formation processes of galaxies and clusters, and their environments \\citep{2007ApJ...671.1135K, 2007ApJ...654...53M, 2008ApJ...681.1076D,2009RAA.....9...41F, 2010MNRAS.404.1137B,2010MNRAS.407..581R,2010arXiv1007.0612W, 2011arXiv1104.1566V}. For example, major mergers of dark matter haloes may play an important role in forming their shapes \\citep{1993ApJ...418..544V}. Mismatches between luminous galaxy shapes and total mass shapes provide evidence for the existence of dark matter \\citep{2010A&A...524A..94S}. Numerical simulations under different assumptions predict dark matter haloes with different shapes \\citep{2002sgdh.conf..109B, 2005ApJ...627..647B}. For instance, collisionless CDM predicts prolate triaxial haloes \\citep{2006MNRAS.367.1781A}, while simulations with gas cooling suggest more spherical shapes of dark matter haloes \\citep{2004ApJ...611L..73K}. Moreover, taking into account the baryonic cooling effect, inner halo shapes are often found to be more spherical than their outer counterparts \\citep{2004IAUS..220..421S,2004ApJ...611L..73K}.  In simulations without baryonic cooling, however, the opposite trend has been seen, i.e.  the outer parts of haloes appear to be rounder than inner parts \\citep{1991ApJ...378..496D}.  The shapes of cluster-sized haloes or galaxy-sized haloes can be studied with different approaches, such as X-ray observations and the Sunyaev-Zeldovich effect \\citep{2000A&A...364..377R, 2004ApJ...601..599L, 2004ApJ...617..847W, 2005ApJ...625..108D, 2006ApJ...645..170S}. One can also use the spatial distribution of satellite galaxies as a tracer for the shapes of their host dark matter haloes \\citep{2008MNRAS.385.1511W, 2010ApJ...718..762W}.  Gravitational lensing provides a powerful tool to study the mass distributions of clusters of galaxies as well as galaxy haloes \\citep[see][for reviews on weak lensing]{2001PhR...340..291B, 2003ARA&A..41..645R,2006glsw.book.....S, 2008PhR...462...67M}. This is because gravitational lensing probes the distribution of matter regardless of whether it is luminous or dark. Weak lensing techniques have been used for cluster mass reconstructions \\citep[e.g.][]{2006ApJ...648L.109C,2008ApJ...687..959B} and for determining halo ellipticities \\citep{2004ApJ...606...67H, 2009MNRAS.393.1235C, 2010ApJ...721..124D, 2010MNRAS.tmp..941H}. The method of stacking galaxies to determine galaxy halo ellipticity was proposed in \\citet{2000ApJ...538L.113N} and has been used to determine ellipticities of both cluster- and galactic-sized haloes \\citep{2009ApJ...695.1446E, 2006MNRAS.370.1008M}. For example, a mean ellipticity of 0.46 has been found from a sample of 25 clusters using Subaru data \\citep{2010MNRAS.405.2215O}.  Flexion has been studied as the derivative of the surface mass density and the shear, responds to small-scale variations in the projected mass distribution \\citep{2002ApJ...564...65G,2005ApJ...619..741G,2006MNRAS.365..414B}. Different techniques have been developed to measure flexion \\citep{ 2006ApJ...645...17I, 2007ApJ...660..995O, 2007MNRAS.380..229M, 2008A&A...485..363S, 2010arXiv1011.3041V,2011arXiv1101.4407F, 2011arXiv1103.0551C}. Due to the complexity of deconvolving for the point spread function, current lensing image measurements are not accurate enough to fully study flexion. It has only been implemented on limited samples of real data. \\citet{2008ApJ...680....1O} and \\citet{2007ApJ...666...51L} have performed flexion measurements to detect substructures in the cluster Abell 1689 using Subaru and HST data. Though flexion noise is still poorly understood, flexion can potentially contribute to cosmology in several aspects, such as exploring dark matter haloes of galaxies and clusters, especially substructures \\citep{2009MNRAS.395.1438L,2010MNRAS.409..389B,2010arXiv1008.3088E}, understanding the large scale structures through cosmic flexion \\citep{2010arXiv1003.5003M,2011arXiv1101.4769S}. \\citet{2011MNRAS.412.1023H} also proposed to reduce the distance measurement errors of standard candles using lensing shear and flexion maps. In \\citet{2009MNRAS.400.1132H}, galaxy halo ellipticity has been studied with flexion. Recently, \\citet[][hereafter ES11]{2011A&A...528A..52E} developed a new approach using the ratio of the tangential to radial flexion to estimate the halo ellipticity as a whole. The approach is independent of lens strength, i.e. the mass or the lens redshift. In particular, the flexion ratio is independent of the halo-centric distance. In this paper, for the first time, we use the flexion ratio to estimate how the halo ellipticity varies with radius.  In reality, there are several complications. For example, the background galaxy number density may not be sufficiently high, and galaxies may not be circularly symmetrically distributed. Even worse the intrinsic and other noises will cause an overestimate for the mass ellipticity. In this paper, we introduce a weight function and a method that combines the first and second order statistics of flexion ratios to reduce the bias (see Appendix A), we test our method with numerical simulations (Sect. 3). Our results are discussed in Sect. 4. The cosmology that we adopt throughout this paper is a $\\Lambda$CDM model with $\\Omega_{\\Lambda}=0.75$, $\\Omega_{m}=0.25$, $\\sigma_8=0.9$ and a Hubble constant $H_0=73\\, {\\rm km}\\,{\\rm s}^{-1}{\\rm Mpc}^{-1}$. ",
        "conclusions": "In this paper, we have studied the flexion ratio method for estimating shapes of elliptical dark matter haloes. We performed numerical simulations to study our ellipticity estimators and obtained the halo ellipticity as a function of radius. In particular, for galactic-sized haloes, we are able to measure the mean variation of ellipticity with radius by stacking several galaxy samples. For cluster haloes, we can estimate the ellipticity as a function of radius for an individual galaxy cluster. We also point out that the minimum $\\robs$ provides another approach for determining the centre of galaxy clusters. Moreover, we introduced a weight function (in Appendix) to suppress the bias due to the asymmetric distribution of background galaxies. The variance of our estimate becomes significantly smaller after employing the weight function. We derived the second order statistics of the flexion ratio $\\langle r^2\\rangle$ and its analytical relation to $\\langle r\\rangle$. By combining $\\langle r\\rangle$ and $\\langle r^2\\rangle$, in principle the bias due to intrinsic noise can be subtracted. We provided a method to reduce the bias based on a linear approximation (For more details, see Eq. \\ref{eq:delta}).  We however notice that there are some difficulties in estimating the halo ellipticity. \\begin{itemize} \\item First of all, the noise model that we employ in this paper includes only intrinsic noise and numerical noise in the simulations. In reality there are other sources of noise, e.g. noises due to the point spread function and pixelisation. Moreover, the noise of flexion measurement appears different using different methods \\citep{2007ApJ...660.1003G}. Thus the noise in practice may be non-Gaussian. Therefore, the results presented here are based on somewhat idealised observations. \\item The difficulty in reducing bias for small ellipticity haloes is mainly due to the lack of knowledge for the flexion noise behaviours. Moreover, the linear approximation that we employed is based on the assumption that the bias is independent of the halo ellipticity and is smaller than the flexion ratio signal. In reality, neither of these assumptions may be valid. In particular, for small elliptical haloes, our measurement will be dominated by intrinsic and other noises. \\item The optimal binning to the data is also difficult since we do not know how the ellipticity changes with radius. A coarse binning can only estimate the ``mean'' ellipticity. Moreover, as we show in the figures, flexion data on different radii have different signal-to-noise ratio (S/N), thus better data at small radii may be polluted by the low S/N data at large radii. On the other hand, if many bins are used, not enough galaxies will be contained within each bin. \\item The degeneracy of substructures and halo ellipticity: substructures behave similarly to intrinsic flexion in our estimate, but at different magnitudes. Substructures are therefore another kind of ``noise''. On the other hand, flexion can be used to measure substructures in galaxies or clusters (e.g. their mass function and spatial distribution). Fig.\\ts\\ref{fig:sigma} is a rough estimate by comparing the $\\sigma_F$ magnitude without any knowledge of substructures or ellipticity. We can see that the combination of different approaches may constrain both shapes and substructures in dark matter haloes. Again good knowledge of flexion noise will be necessary. Also, a realistic model of the substructure mass function, and the spatial distribution of substructure will be helpful. \\end{itemize}  A high number density of background galaxies is desirable.  Current surveys such as COSMOS may have sufficiently high galaxy number densities to enable our method to constrain mass ellipticities on a few radial bins for an individual galaxy cluster. This is again based on a simple noise model. In reality, the noise of high magnitude images is larger. The number density of high quality images using current observations may not as high as we used. For galaxy-sized haloes, we can perform stacking galaxy-galaxy lensing. Notice that different from the shear method, we do not need to align the major axis of the foreground galaxies.  We do, however, assume that we can select dark matter haloes with similar shapes by carefully choosing the foreground galaxies according to their luminosity, colour and shape. The different morphologies will cause uncertainty in our estimate. On the other hand, the number density of galaxy-galaxy pairs will be very large in current and future surveys, i.e. the number of stacked galaxies in reality can be much higher than what we used in our test. Fig.\\ref{fig:stack} indicates we can measure $\\epsilon$ to standard deviation $\\sigma_{\\epsilon}\\approx 0.05$ by stacking 20 foreground galaxies for $\\epsilon\\approx 0.2$. With a much larger number of galaxies available in future observations, the accuracy of this measurement can be improved. The improvement will be Poissonian if there are no other sources of systematic errors. An accurate measurement of how the ellipticity varies with radius will provide a strong test of galaxy formation models."
    },
    "1107/1107.2651_arXiv.txt": {
        "abstract": "Most massive galaxies contain supermassive black holes (SMBHs) in their cores. When galaxies merge, gas is driven to nuclear regions and can accrete onto the central black hole. Thus, one expects to see dual active galactic nuclei (AGNs) in a fraction of galaxy mergers. Candidates for galaxies containing dual AGNs have been identified by the presence of double-peaked narrow [O III] emission lines and by high spatial resolution images of close galaxy pairs.  Spatially resolved spectroscopy is needed to confirm these galaxy pairs as systems with spatially separated double SMBHs.  With the Keck 2 Laser Guide Star Adaptive Optics system and the OH Suppressing InfraRed Imaging Spectrograph near-infrared integral field spectrograph, we obtained spatially resolved spectra for SDSS J09527.62+255257.2, a radio-quiet quasar shown by previous imaging to consist of a galaxy and its close (1\\farcs0) companion. We find that the main galaxy is a Type 1 AGN with both broad and narrow AGN emission lines in its spectrum, while the companion galaxy is a Type 2 AGN with narrow emission lines only. The two AGNs are separated by 4.8 kpc, and their redshifts correspond to those of the double peaks of the [O III] emission line seen in the Sloan Digital Sky Survey spectrum.  Line diagnostics indicate that both components of the double emission lines are due to AGN photoionization.  These results confirm that J0952+2552 contains two spatially separated AGNs.  As one of the few confirmed dual AGNs at an intermediate separation of \\textless 10 kpc, this system offers a unique opportunity to study galaxy mergers and their effect on black hole growth. ",
        "introduction": "%  In $\\Lambda$CDM cosmology, galaxies regularly interact and merge; hierarchical galaxy formation and evolution dictates that galaxies undergo multiple mergers as they evolve \\citep[e.g.,][]{blumenthal1984, spergel2007}.  We now know that most massive galaxies harbor supermassive black holes (SMBHs) in their centers \\citep{richstone1998, kormendy1995}. One picture explaining active galactic nuclei (AGNs) invokes mergers of gas-rich galaxies; the gravitational perturbations due to the merger allow gas to flow into the galactic center and trigger accretion onto the black hole \\citep{hernquist1989, kauffmann2000, hopkins2008}.  A complication in this picture is the AGN's duty cycle, or how long each AGN remains active during the merger.  If AGN activity is triggered by galaxy mergers, AGN pairs should be observable in at least some galaxy mergers.  Thus, the existence and statistics of dual AGNs provide an important probe into hierarchical galaxy formation models, accretion-triggering mechanisms, galaxy merger rates, and SMBH growth \\citep[and references therein]{yu2011arxiv}.  AGN pairs have been found and resolved both spatially and spectrally at a wide range of separations.  Hundreds of AGN pairs are known at \\textgreater 10 kpc separations \\citep[e.g.,][]{myers2007, myers2008, hennawi2010, green2010, piconcelli2010}, while only three candidate binary AGNs are known at \\textless 10 pc separations \\citep{rodriguez2006, boroson2009, decarli2010}.  Only five spatially and spectrally confirmed dual AGNs are known with intermediate separations of between 0.1 and 10 kpc: LBQS 0103-2753 \\citep{junkkarinen2001}; NGC 6240 \\citep{komossa2003}; Arp 299 \\citep{ballo2004}; J1420+5259 \\citep{gerke2007}; Mrk 463 \\citep{bianchi2008}; Mrk 739 \\citep{koss2011}; and J1027+1749 \\citep{liu2011}. Eighteen candidate AGN pairs with spatially unresolved double-peaked emission lines and resolved double spatial structure have been observed by \\citet{fu2011} and \\citet{ rosario2011}.  Potential dual AGNs have been selected using the following two methods. \\begin{enumerate} \\item Observing velocity offsets between emission lines, such as \\oiii, potentially from two AGNs \\citep{smith2010, zhou2004, wang2009, liu2010a, liu2010b, comerford2009a, comerford2009b, XuKomossa2009, gerke2007}. \\item Imaging multiple potential AGNs in or close to a single host galaxy \\citep{junkkarinen2001, komossa2003, comerford2009b, fu2011, rosario2011, liu2011}. \\end{enumerate} To \\underline{confirm} a potential dual AGN, spatially resolved spectroscopy is needed to prove that each resolved source has a unique AGN spectrum.  Without spatially resolved spectroscopy to match the observed galaxies to their corresponding double-peaked emission lines, the double peaks of spectral lines may be due to a chance superposition of two objects, a recoiling SMBH \\citep{bonning2007, civano2010, guedes2011},  jets interacting with the surrounding medium \\citep{rosario2010}, outflows from a single AGN \\citep{fischer2011}, or rings of star formation.   Similarly, without spatially resolved spectroscopy, the multiple bright cores imaged in or around one galaxy may be a chance superposition, a recoiling SMBH, jets interacting with the surrounding medium, outflows from a single AGN, gravitationally lensed sources \\citep[and references therein]{hennawi2006}, or starbursts.  High-energy X-ray observations are another unambiguous way to confirm true dual AGNs.  In the next section, we discuss W. M. Keck Observatory adaptive optics imaging of \\obj, a potential pair of bright AGNs with double-peaked emission lines in the Sloan Digital Sky Survey (SDSS) spectrum. In Section 3, we describe our spectroscopic observations of \\obj~using adaptive optics spectroscopy from Keck's OH Suppressing InfraRed Spectrograph (OSIRIS) and the resulting data reduction.  In Section 4, we match resolved spatial components seen by NIRC2 and OSIRIS to the SDSS double-peaked emission lines, and discuss what type of physical mechanisms are generating these emission lines.  Section 5 presents our conclusions.   We adopt a concordance cosmology with $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{\\Lambda} = 0.7$.  All quoted wavelengths are in vacuum units. ",
        "conclusions": "To confirm candidate dual AGNs selected via double-peaked \\oiii~emission lines or multiple bright cores imaged in a single galaxy, spatially resolved spectroscopy is required.   For candidate dual AGNs with larger separations between components (\\textgreater 1\\farcs5), long slit spectrographs such as the Low Resolution Imaging Spectrometer can be used. For candidates with small separations (\\textless1\\farcs5), AO systems are best used to obtain spatially resolved spectra which clearly separated the contributions from the two galaxies.  We used the Keck 2 LGS AO system and the OSIRIS near-infrared integral field spectrograph to obtain spatially resolved spectra for \\obj, a radio-quiet system with double-peaked \\oiii~lines and two galaxies observed using Keck 2's LGS AO system and NIRC2.  The infrared redshifts of the two AGNs, separated by 4.8 kpc or 1\\farcs0, match well with the redshifts of the SDSS \\oiii~double peaks.  This directly links the double-peaked emission lines to the narrow-line regions of the two observed AGNs.  Line diagnostics indicate that both the companion and the main galaxy are bright due to AGN photoionization rather than through star formation or supernovae.  Spatially resolved spectroscopic observations of further dual AGNs candidates are needed to construct a statistically-significant sample of true dual AGNs.  With such a sample, topics such as the merger rate or the AGN duty cycle can be addressed through comparisons with simulations of multiple mergers such as \\citet{yu2011arxiv,lotza2010,lotzb2010}, and \\citet{lotz2011arxiv}."
    },
    "1107/1107.0979_arXiv.txt": {
        "abstract": "The current state of the art in the modeling of pulsar magnetospheres invokes either the vacuum or force-free limits for the magnetospheric plasma. Neither of these limits can simultaneously account for both the plasma currents and the accelerating electric fields that are needed to explain the morphology and spectra of high-energy emission from pulsars. To better understand the structure of such magnetospheres, we combine accelerating fields and force-free solutions by considering models of magnetospheres filled with resistive plasma.  We formulate Ohm's Law in the minimal velocity fluid frame and construct a family of resistive solutions that smoothly bridges the gap between the vacuum and the force-free magnetosphere solutions. The spin-down luminosity, open field line potential drop, and the fraction of open field lines all transition between the vacuum and force-free values as the plasma conductivity varies from zero to infinity. For fixed inclination angle, we find that the spin-down luminosity depends linearly on the open field line potential drop.  We consider the implications of our resistive solutions for the spin down of intermittent pulsars and sub-pulse drift phenomena in radio pulsars. ",
        "introduction": "\\label{sec:intro}  Pulsar magnetospheres are filled with plasma, and the presence of plasma affects all magnetospheric properties. Although this fact has been well appreciated since the early days of pulsar research, the ability to quantitatively model plasma effects has emerged mainly in the last decade.  Generally speaking, the models of pulsar magnetospheres can be classified according to the amount of plasma supply they assume. At one extreme is the vacuum magnetosphere, which is modeled with the magnetic field of a spinning inclined dipole in vacuum \\citep{Deutsch55}.  As this solution has no plasma, it has no possibility of producing any pulsar-like emission. However, the fact that this field is described by an analytic formula has made it the most widely used framework for calculating pulsar properties. For example, the characteristic spin-down energy loss which yielded the fiducial pulsar field strength of $10^{12}$G or the polarization sweep of the rotating-vector model \\citep{RadCooke69} are guided by the vacuum field shape. Slot-gap and outer gap models of gamma-ray emission from pulsars are also based on this field \\citep{CHR86a, DHR04}.  The next order of approximation is models that allow plasma emission from the surface of the star. These include the original charge-separated model of \\citet[hereafter GJ]{GJ69}, and the space charge-limited models with pair formation \\citep[e.g.,][]{rudermansutherland75,AS79}. These models envision both the regions where plasma shorts out the accelerating electric fields and the vacuum-like regions where acceleration is present (so-called ``gaps''). This approach allows for more realism in studying plasma creation and acceleration, but at the price of being decoupled from the global structure of the magnetosphere, as the modification of the vacuum field due to plasma currents is typically not included. Finally, the models that concentrate on the global magnetospheric properties assume that abundant plasma exists throughout the magnetosphere and in the wind. This plasma shorts out the accelerating electric fields and provides the corotation of field lines with the star. Such models include the relativistic magnetohydrodynamic (MHD) description of the magnetosphere and its limit for low-inertia magnetically-dominated plasmas, the ``force-free'' description. Which of these different regimes is applicable to real pulsars may ultimately depend on the physics of plasma supply in the magnetosphere.  Currently, quantitative solutions of the global magnetospheric structure exist only for the vacuum limit \\citep{Deutsch55} and for the limit of abundant plasma in force-free electrodynamics (see \\citealt{CKF99, Gruzinov05, Timokhin06, McKinney06, Komissarov06}\\footnote{We note that \\cite{Komissarov06} has also calculated the structure of the aligned rotator in the relativistic MHD limit including particle inertia.}  for aligned rotators, and \\citealt[hereafter S06]{Spitkovsky06}; \\citealt{CK10} for pulsars with arbitrary inclinations).  The real pulsar magnetosphere is likely operating somewhere in between these limits, with various accelerating gaps, regions of pair production, and strong current sheets likely causing local violations of the ideal MHD constraint, $\\vec{E}\\cdot \\vec{B}=0$. Knowing the structure of the magnetosphere, including such non-ideal effects, would be very useful for calculating the properties of pulsar emission. Indeed, currently the ideal force-free models that include the back-reaction of plasma currents on the field structure lack any accelerating fields by construction, and thus cannot be used to directly predict the spectra of gamma-ray radiation observed by Fermi GST.  One way to reintroduce accelerating electric fields in the magnetosphere is to allow finite resistivity of the plasma. Several formulations of resistive force-free equations have been proposed, most notably by \\cite{Lyutikov03} and \\cite{Gruzinov07,Gruzinov08}. In resistive MHD, the Ohm's law can be unambiguously defined in the proper frame of the fluid by relating the current in that frame to the electric field through $\\vec{j}_{\\rm fluid}=\\sigma \\vec{E}_{\\rm fluid}$, where $\\sigma$ is plasma conductivity \\citep[see, e.g.,][]{Lichnerowicz67, Palenzuela09}. In the force-free system, however, the fluid velocity along the magnetic field is unknown, and only the transverse velocity can be obtained from the electromagnetic fields. This introduces some freedom in prescribing the Ohm's law in the force-free picture. In the prescription proposed by \\cite{Lyutikov03}, the parallel velocity along the field was taken to be zero. In ``Strong-Field Electrodynamics'' \\citep[hereafter SFE]{Gruzinov07,Gruzinov08}, the Ohm's law was formulated in the frame that moves along the field lines with such speed that the charge density in that frame vanishes. Despite the fact that such a frame formally exists only in space-like regions, SFE prescription appears to give smooth numerical solutions throughout the magnetosphere \\citep{Gruzinov11}, and, most importantly, does not explode in current sheets where magnetic fields can go through zero. Both formulations are arbitrary, however, because the real fluid velocity does not have to follow either frame assumption. As an additional constraint, it is useful, therefore, to construct a resistive formulation that reproduces physical solutions expected at the extremes of very large and very small conductivity of the plasma, namely the force-free limit ($\\sigma\\to\\infty$) and the limit of vacuum electromagnetism ($\\sigma\\to 0$). Below we describe a resistive prescription that generalizes these schemes and combines the correct limiting behavior of Lyutikov's scheme with current-sheet stability of Gruzinov's formulation. We then apply this prescription to numerically calculate the structure of resistive magnetospheres in pulsars.  In the resistive force-free picture of the pulsar magnetosphere, the magnetized neutron star is thought to be surrounded by an abundant massless plasma with finite conductivity, so that not all accelerating fields are shorted out. For simplicity and as a proof of principle, we only consider the unrealistic case of constant conductivity throughout space. More complicated prescriptions will be studied elsewhere. Our finding is that using our formulation of the resistive force-free electrodynamics we can construct a family of magnetospheres that smoothly transition from the Deutsch vacuum solution to the ideal force-free magnetosphere as the conductivity of the plasma is increased. Such intermediate magnetospheres possess interesting properties that we discuss in this paper.  We study the variation of the spin-down power with magnetic inclination angle as a function of plasma conductivity and relate it to physical conditions such as the effective potential drop on the open field lines. In \\S\\ref{sec:currentder} we discuss the derivation of the resistive force-free prescription, in \\S\\ref{sec:simulations} we describe our code for solving the equations and present sample magnetospheric solutions. Discussion and potential applications to pulsar physics are in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have presented a continuous one-parameter family of pulsar magnetosphere solutions that span the range between the vacuum and force-free limits.  Each solution of the family is characterized by the value of the conductivity parameter, $\\sigma/\\Omega$, which is related to the maximum potential drop that a test particle can experience as it moves along field lines, $V_{\\rm drop}$.  The zero conductivity limit, $\\sigma/\\Omega\\to0$, $V_{\\rm drop}\\to V_{\\rm drop,vac}$, yields the vacuum magnetosphere.  It shows a substantial effective potential drop that makes up a significant fraction of the rotation-induced potential difference between the pole and the equator of the star ($20{-}60\\%$ depending on the inclination).  The high conductivity limit, $\\sigma/\\Omega\\to\\infty$, yields essentially an ideal force-free magnetosphere. Unlike in vacuum, abundant magnetospheric charges short out any potential differences along field lines, leading to vanishing effective potential drop, $V_{\\rm drop,ff} \\to 0$.  While in the vacuum case nearly all field lines, even those that extend beyond the light cylinder, eventually return to the star and are formally closed, in the ideal force-free case a fraction of field lines open up and reach infinity. Our simulations show that resistive high-$\\sigma$ pulsars spin down at least $3$ times faster than resistive low-$\\sigma$ and vacuum pulsars (for the same value of inclination), in agreement with earlier ideal force-free simulations (S06).  Our resistive solutions bridge the gap between the force-free and vacuum limits by having intermediate values of major magnetospheric parameters, such as effective potential drop, $0<V_{\\rm drop}<V_{\\rm drop,vac}$, the fraction of open field lines, and the spin-down rate.  Before discussing possible observational implications of these solutions, we note that our approach to modeling resistive pulsar magnetospheres is quite simplistic: we use a form of Ohm's Law in which we neglect several terms \\citep{Meier04} and we assume constant conductivity $\\sigma/\\Omega$ throughout all space.  The finite bulk conductivity used in this work can be thought of as not due to collisions, but due to the insufficient plasma supply in the magnetosphere that cannot short out all the accelerating fields.  It is possible that real magnetospheres have a lower (anomalous) conductivity in the current layer and the current sheet than in the rest of the magnetosphere. We might also expect a lower value of conductivity on the open field lines than on the closed field lines. Since our models do not explicitly model the plasma fluid (\\S\\ref{sec:intro}), the plasma velocity component along the direction of magnetic field is unconstrained, which necessitates the choice of a frame in which to write the Ohm's Law.  We choose the so-called minimum velocity frame, which has a number of attractive properties (see \\S\\ref{sec:currentder}), yet this is not a covariant choice. With these important caveats in mind, let us now for the sake of argument take our model seriously as a description of pulsar magnetosphere and consider the consequences.  Intermittent pulsars offer a unique test bed of pulsar theory (e.g., PSR B1931+24, \\citealt{Kramer06}, and PSR J1832+0029, \\citealt{Lyne09}; see also \\citealt{Wang07}, \\citealt{Zhang07}, and \\citealt{Timokhin10}).  Such pulsars switch between an ``on'', radio-loud, state in which they behave like normal radio pulsars, and an ``off'', radio-quiet, state in which they produce no detectable radio emission.  The two intermittent pulsars, for which published data exist, have quite different duty cycles: PSR B1931+24, with a period $P\\approx0.8$~s, cycles through the ``on''--''off'' sequence of states approximately once a month, whereas PSR J1832+0029, with a period $P\\approx0.5$~s, kept quiet for nearly two years.  The spin-down rate for each of these pulsars is larger in the ``on'' state than in the ``off'' state by a factor $f_{{\\rm on}\\to{\\rm off}}\\simeq1.5$.  Such a substantial difference in spin-down rates suggests that the pulsar magnetosphere undergoes a dramatic reconfiguration as it transitions between the ``on'' and ``off'' states, yet such a transition was reported for PSR B1931+24 to take place in just over $10$ pulsar periods.  \\citet{Kramer06} propose that in the ``on'' state, plasma fills the pulsar magnetosphere and supports (poorly understood) plasma processes that produce radio emission.  In this picture, the pulsar transitions to the ``off'' state when, due to some unknown trigger, magnetospheric pair production shuts off: the remaining plasma leaks off the open field lines and the pulsar goes radio-quiet.  In this scenario, we expect the pulsar in the ``on'' state to spin down due to force-free energy losses and in the ``off'' state due to vacuum dipole losses. However, this picture is ruled out by observations: the force-free spin-down rate is larger than the vacuum spin-down rate by a factor, $f_{{\\rm ff}\\to{\\rm vac}}\\ge3$ (S06), that is clearly incompatible with the observed value, $f_{{\\rm on}\\to{\\rm off}}\\simeq1.5$ \\citep[e.g.,][]{BN07}.  Our resistive magnetospheres provide a possible resolution to this problem since they have intermediate spin-down rates between vacuum and force-free.  Let us associate the ``on'' state with the force-free magnetosphere and the ``off'' state with one of our intermediate resistive magnetospheres.  According to equation (\\ref{eq:spindown}), for the ratio of spin-down powers in these two states to equal $f_{{\\rm on}\\to{\\rm off}}$, the potential drop has to be ${V_{\\rm drop}} = V_0{(f_{{\\rm on}\\to{\\rm off}}-1)\\left(0.9+1.1 \\sin^2\\alpha\\right)}/{4.5 f_{{\\rm on}\\to{\\rm off}}}$.  For $f_{{\\rm on}\\to{\\rm off}}\\approx1.5$, the ratio $V_{\\rm drop}/V_0$ increases monotonically with increasing inclination angle, $\\alpha$, from $0.07$, for an aligned rotator ($\\alpha=0^\\circ$), to $0.15$, for an orthogonal rotator ($\\alpha=90^\\circ$). Plugging in for $V_0$, using the measured spin-down parameters of PSR B1931+24 (we associate the observed spin-down rate in the ``on'' state with the spin-down rate of the force-free magnetosphere), and assuming a neutron star mass, $M_*=1.4 M_\\odot$, we obtain $V_{\\rm drop} = 1.5\\times10^{16}(0.9+1.1\\sin^2\\alpha)^{1/2}~[{\\rm V}]$.  This value significantly exceeds the characteristic value of the polar cap potential drop expected for pulsars \\citep{GJ69, Arons09}, so the pulsar should have no difficulty in creating pairs and shorting out the large potential. The reason for this discrepancy is our simplifying assumption that the conductivity is constant throughout the magnetosphere.  This assumption causes the potential drop not just across the polar cap but across the whole stellar surface.  In real pulsars, we expect that the closed field line region is filled with plasma and is highly conductive.  This plasma shorts out parallel electric field in the closed field line region and screens the potential drop everywhere except across the polar cap.  The potential drop across the polar cap, $V_{\\rm pc}$, is $\\approx(R_*/R_{\\rm LC})$ times the pole-to-equator potential drop, $V_0$, which gives \\begin{align} V_{\\rm drop} &\\sim V_{\\rm pc}\\frac{(f_{{\\rm on}\\to{\\rm off}}-1)}{4.5 f_{{\\rm on}\\to{\\rm off}}}\\\\ &=3.8\\times10^{12}~[{\\rm V}], \\notag \\end{align} for PSR B1931+24, where we assumed $R_*=10$~km.   Addition of resistivity can, in principle, allow the detailed study of pulsar spin-down and the determination of pulsar braking indices.  For a spin-down law $\\Omega = \\Omega(t)$, the braking index is defined as $ n \\equiv {\\Omega \\ddot\\Omega}/{\\dot\\Omega^2} $.  In order to measure a braking index numerically, we run a series of simulations for different values of angular frequency, $\\Omega$, and map out the dependence, $L(\\Omega)$.  We find that $L(\\Omega)\\propto \\Omega^4$ for both purely force-free and vacuum solutions, which translates into a value of the braking index, $n=3$, as we will see below.  This is larger than the range of observed values, $n\\simeq2.5{-}2.8$ \\citep{Livingstone}.  One way to lower the braking index is to allow the evolution of the extent of the magnetosphere (the radius of the Y-point) to lag behind the outward expansion of the light cylinder due to spin down \\cite[e.g.,][]{ContopSpitkovsky06}. The extent of the Y-point is controlled by the rate of reconnection at the edge of the magnetosphere. We calculated the braking index in our simulations by comparing the magnetosphere with constant conductivity spun up to different periods, and we do not find significant deviations from $n=3$.  This rapid spinup biases our answer, however, because the Y-point extends to the light cylinder within one rotation period.  It is not entirely clear that we would obtain the same braking index if we allowed the pulsar period to increase during a single simulation as the pulsar spins down and the light cylinder slowly recedes.  The extent of the Y-point may depend on the past history of the pulsar. To do this calculation, we need to increase the pulsar period several times within a single simulation and measure the resulting spin-down luminosity after each increase.  We must also ensure that the reconnection at the Y-point is controlled by the resistivity we impose and not by numerical effects.  A more accurate braking index calculation would require much longer simulations, and possibly higher spatial resolution near the Y-region.  Another possibility that can affect the braking index is the evolution of conductivity with time. Such time-evolution translates into evolution in pulsar dimensionless luminosity, $\\ell = L/L_0$ (see Fig.~\\ref{dipole}), and leads to $n\\neq3$: \\begin{equation} \\label{eq:n} n = \\frac{\\Omega \\ddot\\Omega}{\\dot\\Omega^2} = 3 - 2 \\frac{\\mathrm d\\ell/\\mathrm dt}{\\ell/\\tau}, \\end{equation} where $\\tau=P/2\\dot P$ is the pulsar age, and we used the scaling, $L_0\\propto \\Omega^4$. If dimensionless plasma conductivity does not change in time, $\\sigma/\\Omega=$ constant, then $\\mathrm \\ell=$ constant and equation (\\ref{eq:n}) gives $n = 3$.  Purely force-free and purely vacuum pulsars, discussed above, fall into this category.  As a pulsar ages, it is natural to expect that its plasma supply depletes and the pulsar becomes progressively more vacuum-like: in other words, both $\\sigma/\\Omega$ and $\\ell$ decrease in time (see Fig.~\\ref{dipole}).  For such pulsars, $\\mathrm d\\ell/\\mathrm dt<0$, and equation (\\ref{eq:n}) gives $n > 3$.  Of course, the converse is also true, and the braking index lower than $3$ would result if a pulsar becomes progressively more force-free--like, i.e., if pulsar dimensionless luminosity increases in time, $\\mathrm d\\ell/\\mathrm dt>0$. Presently, however, it is not clear why such behavior would be physically expected. Thus, the addition of bulk resistivity does not seem to easily solve the braking index puzzle, and a more elaborate explanation is still required.  Our resistive magnetospheres naturally lend themselves to modeling the sub-pulse drift phenomena observed in a number of pulsars \\cite[see e.g.,][]{Askegar01,DeshpandeRankin,Rankin03,ContopSpitkovsky06}.  The following discussion takes place in the corotating frame, where we define electromagnetic fields by equations (\\ref{eq:corotEfield}) and (\\ref{eq:corotBfield}).  In this frame there is transverse particle drift with velocity of order $v_{\\rm drift}\\sim \\vec{E}_{\\perp}\\times \\vec{B}/B^2$.  This drift velocity can be thought of as the minimal velocity of particles, ignoring motion along the magnetic field lines. The transverse electric field $\\vec{E}_{\\perp}$ reverses sign through the axis of maximum potential drop, and so there is rotation of field lines about this axis.  In general, this axis is not in the direction of the magnetic moment, and in fact does not even need to be a straight line.  If we consider a simplified picture of the magnetosphere in which there is no parallel electric field in the closed field line region, the potential drop across open field lines is given roughly by the longitudinal potential drop along field lines. The differential rotation can then be roughly related to the longitudinal potential drop by \\begin{equation} \\Delta\\Omega = \\Delta V_{\\rm drop}/(B_p R_\\perp^2) = \\Delta V_{\\rm drop} / \\Psi_{\\rm cap}. \\label{eq:domegapc} \\end{equation}  In the laboratory frame, the field lines are undergoing rotation about the pulsar spin axis (with the angular frequency $\\Omega$), and in addition they are rotating differentially around the maximum potential axis in a retrograde sense.  The differential rotation of plasma can explain sub-pulse emission features that drift with respect to the otherwise periodic light curve of the pulsar.  Our result is in contrast to the standard picture in which the plasma precesses about the magnetic axis \\citep{rudermansutherland75}.  The magnetic colatitude parameter in the cartographic transformation of, e.g. \\cite{DeshpandeRankin}, must then be modified accordingly.  To better visualize the differential rotation in the context of our resistive magnetospheres, consider the full drift velocity \\begin{equation} v_{\\rm drift}=\\frac{\\vec{E}\\times\\vec{B}}{B^2+E_0^2}c, \\end{equation} again in the corotating frame.  Fig.~\\ref{velocity} shows magnetic field lines in the $\\vec{\\mu}-\\vec{\\Omega}$ plane for a $60^{\\circ}$ inclined dipole at conductivity $(\\sigma/\\Omega)^2=4$.  Color indicates out-of-plane drift velocity, with red (blue) representing velocity into (out of) the page.  The velocity has been rescaled by raising its magnitude to the $1/2$ power.  The axis of maximum potential extends from the star roughly halfway between the rotation and magnetic axes.  The differential rotation of plasma about this axis is denoted by the transition from red to blue across the axis, indicating a reversal in the sign of the out-of-plane drift velocity. Note that the axis bends off in the direction of open field lines beyond the light cylinder.  These corotating frame drift velocities can be quite substantial.  The maximum illustrated drift velocity occurs near the current sheets and corresponds to a velocity of $0.8c$.  The differential rotation is much stronger than that implied by sub-pulse drift, but we expect the qualitative features of differential rotation to persist in solutions with more realistic open field line potential drops.  At lower conductivities the magnitude of differential rotation about the axis of maximum potential drop increases, and the conductivity is a parameter that can be tuned to attempt to match observed sub-pulse emission features.  Incidentally, the bundle of field lines surrounding the axis of the maximum potential corresponds to the field lines that carry the largest distributed current density on the polar cap (see Fig. 4 in \\citealt{BS10}). If the strength of the current is associated with radio emission, then the core emission would have a centroid that is offset from the magnetic pole. The offset could be as large as half of the polar cap radius. This would introduce potentially significant modifications to the polarization sweep of the core radio emission and cause deviations from the expected S-curve of \\cite{RadCooke69}. For a number of pulsars, it would imply differences in the inclination and viewing angles inferred from the shape of the polarization sweep.  \\begin{figure}[t] \\centering \\includegraphics[scale=.9]{f5} \\caption{Magnetic field lines in the $\\vec{\\mu}-\\vec{\\Omega}$ plane for a $60^{\\circ}$ inclined dipole at conductivity $(\\sigma/\\Omega)^2=4$.  Color represents the out-of-plane drift velocity in the corotating frame, with blue color representing the out of the page direction.  There is retrograde differential rotation about the axis with maximum potential drop, situated roughly halfway between the rotation and magnetic axes when inside the light cylinder, as well as about the boundary of the closed field line region.  The maximum value on the color table corresponds to a drift velocity of $0.8c$.}\\label{velocity} \\end{figure}  We conclude with prospects for future research.  We will generalize our assumption that $\\sigma/\\Omega$ is constant in the magnetosphere. In addition to having higher conductivity in the closed zone, we will experiment with prescriptions for anomalous resistivity in the current sheet.  These improvements will yield a more realistic magnetospheric structure, that can be accurate enough for geometrical modeling of gamma-ray light curves from pulsars. The gamma-ray pulse formation is sensitive to the geometry of the magnetosphere, and, in particular, to the field lines near the current sheet \\citep{BS10}. Deviations from the force-free geometry may be important for modeling the light curves of older pulsars which require wider gaps in the outer-gap models \\citep{Watters09}.  AT acknowledges support by the Princeton Center for Theoretical Science fellowship and by the National Science Foundation through TeraGrid resources provided by NICS Kraken under grant number TG-AST100040. AS is supported by NSF grant AST-0807381 and NASA grants NNX09AT95G and NNX10A039G.  We thank the referee, Ioannis Contopoulos, for comments that helped improve this paper.  The simulations presented in this paper were performed on computational resources supported by the PICSciE-OIT High Performance Computing Center and Visualization Laboratory. This research used resources of the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the US Department of Energy under contract No. DE-AC02-05CH11231."
    },
    "1107/1107.0148_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  Despite the impact on their surroundings, the formation and early evolution of massive stars is poorly constrained, primarily because of their scarcity and short lifetimes.  Most OB stars form in stellar clusters and associations, and are still partly hidden in their natal molecular cloud.  To detect the stellar content of these regions directly, near-infrared observations, especially spectroscopy, have proven to be a powerful method to find and characterize the newborn OB stars.  A pure photometric characterization of young stellar clusters is strongly hampered by highly varying extinction, unknown distances and infrared excess of the young cluster members.   A spectroscopic identification, however,  results in a unambiguous identification of the (massive) stellar content. Stellar properties, like effective temperature and luminosity, are derived based on the spectral features and extinction, distance and infrared excess can be determined reliably \\citep[e.g.][]{Blum00, Hanson02,Ostarspec05,Puga06,Bik10,Puga10}. By comparing the effective temperature and  luminosity to main sequence and pre-main-sequence isochrones, the age of a stellar cluster can be derived more reliably  than from photometry alone \\citep{Bik10}.  When studying larger samples of massive stars in the same star cluster, estimates  of an age spread in the cluster can be obtained.  Based on the massive stars, \\citet{Clark05} finds very little age spread (less than 1 Myr) in starburst clusters. However, in Orion an age spread has been found based on the PMS stars \\citep{daRioOrion10}. Additional evidence for an age spread has been found in S255 where the most massive star is still deeply embedded and driving an outflow, while the PMS population has an age of 1-3 Myr \\citep{Wang11}.  We present deep LUCI near-infrared JHKs imaging as well as K-band multi-object-spectroscopy of the massive stellar content of W3 Main. This allows for the first time a spectral classification of the massive stellar content. Using their derived spectral type we discuss the evolutionary status of the massive stars as well as the HII regions in  detail. The different evolutionary phases of the HII regions make W3 Main an ideal target to study age spread and the evolution of circumstellar disks around massive stars. ",
        "conclusions": "W3 Main harbors several different evolutionary stages of HII regions, ranging from very young hyper-compact HII (HCHII) regions (few 10$^3$ years), ultra-compact HII (UCHII) regions \\citep[$\\sim 10^5$ years][]{WoodIRAS89} to evolved,  diffuse HII regions (few 10$^6$ years). All these regions are most likely formed out of the  same molecular cloud. This provides the possibility to study the evolution of young HII regions and their stellar content in great detail.  \\citet{Tieftrunk97} derived an evolutionary sequence for the HII regions in W3 Main, based on the morphology of the radio sources. The youngest are the HCHII regions W3 M and W3 Ca, with the UCHII regions W3 F, W3 C and W3 E, slightly older, the compact HII regions W3 B and W3 A even more evolved, and the diffuse HII regions, W3 K and W3 J being the oldest HII regions in W3 Main. This classification can be compared to the ages of the massive stars deduced from their position in the HRD.  We have detected OB stars in three diffuse HII regions, two compact HII regions and three UCHII regions. The HCHII region W3 M harbors the high-mass protostar IRS5 (see Sect. 4.1).  Additionally three stars have no detectable HII region associated with them (IRS N5 and the massive YSOs IRS N7 and IRS N8).  The position of IRS2 in the HRD suggests an age of  2-3 Myr, consistent with its location in a relatively evolved compact HII region (W3 A) with a similar or younger age derived for  IRS3a,  harbored in a younger UCHII region. For the lower-mass stars (late O, early B) the isochrones are too close to each other to derive any age information. For IRS N1, located inside the UCHII region W3 E, the offset from the main sequence could be explained by high-mass pre-main-sequence evolution \\citep{Hosokawa09}, consistent with the expected young age of the UCHII region.  A similar sequence can be seen in the extinction towards the different HII regions. The extinction varies from A$_{Ks}$ = 0.9 mag for the diffuse HII regions W3 J and W3 K to very high extinction (A$_{Ks}$  = 5.9 mag) towards  W3 F.  Such a sequence in extinction indicates that the stars in the diffuse HII regions have already cleared out their surroundings and destroyed the molecular cloud, while those in UCHII regions are still  embedded in their parental molecular cloud.   The evolutionary sequence of the HII regions seen in the radio morphology, is consistent with an increase in extinction from the older (diffuse HII) regions to the younger (UCHII) regions.  We have detected the photospheres of OB stars from the more evolved diffuse HII region to the much younger UCHII regions, suggesting that the OB stars have finished their formation and cleared away their possible circumstellar disks very fast. Only in the HCHII phase (IRS 5), the massive star is still surrounded by circumstellar material.    Based on the presence of different evolutionary stages of HII regions as well as the location of the most massive stars in the HRD, we can conclude that an age spread of  2-3 Myr is most likely present for the massive stars in W3 Main. A growing number of  young stellar clusters show evidence for an age spread, usually based on the analysis of their PMS population in the HRD. In Orion an age spread of a few Myr has been found  \\citep{Palla99,daRioOrion10}. In starburst clusters, however, upper limits on the age spread of less than 1 Myr have been found for Westerlund 1 \\citep{Clark05} based on its massive stars. This suggests that W3 Main is not formed in one star formation burst as expected for starburst cluster, but, more likely, through a temporal sequence of star formation events.    \\small  %"
    },
    "1107/1107.0238_arXiv.txt": {
        "abstract": "At large magnetic Reynolds numbers, magnetic helicity evolution plays an important role in astrophysical large-scale dynamos. The recognition of this fact led to the development of the dynamical $\\alpha$ quenching formalism, which predicts catastrophically low mean fields in open systems. Here we show that in oscillatory \\aO dynamos this formalism predicts an unphysical magnetic helicity transfer between scales. An alternative technique is proposed where this artifact is removed by using the evolution equation for the magnetic helicity of the total field in the shearing--advective gauge. In the traditional dynamical $\\alpha$ quenching formalism, this can be described by an additional magnetic helicity flux of small-scale fields that does not appear in homogeneous $\\alpha^2$ dynamos. In \\aO dynamos, the alternative formalism is shown to lead to larger saturation fields than what has been obtained in some earlier models with the traditional formalism. We have compared the predictions of the two formalisms to results of direct numerical simulations, finding that the alternative formulation provides a better fit.  This suggests that worries about catastrophic dynamo behavior in the limit of large magnetic Reynolds number are unfounded. ",
        "introduction": "While the possibility, and indeed need, for astrophysical dynamos was recognized quite early \\citep{Lar19}, the study of dynamos has since been troubled by a number of problems. Cowling's anti-dynamo theorem \\citep{Cowling} initially appeared to demonstrate that the entire concept was impossible, though \\cite{Par55} eventually discovered the physics behind what has come to be called the $\\alpha$ effect. Cowling's anti-dynamo theorem was finally shown to be largely inapplicable by analytically solvable dynamos such as the Herzenberg dynamo \\citep{Herzenberg}.  Once the possibility of dynamo action was demonstrated, the development of mean-field $\\alpha$ dynamo theory followed \\citep{Steenbeck}, which describes the generation of poloidal field from toroidal fields.  While the generation of toroidal magnetic fields from sheared poloidal fields is straightforward through the $\\Omega$ effect, the reverse process is tricky.  Without it however, dynamo action is impossible. The $\\alpha$ effect, which relies on helicity (twist) in the fluid motion, allows for the generation of strong large-scale magnetic fields such as those observed in the Universe. It can drive dynamo action on its own (\\alp systems), but as shear is ubiquitous in astrophysics, shear-amplified dynamo action is generally expected to outperform \\alp dynamos. Accordingly, \\aO dynamos, which combine the effects, are expected to be the dominant type of natural astrophysical dynamo \\citep{Shear}.  More recently however, there were indications, first suggested by \\cite{Vainshtein}, that the $\\alpha$ effect decreases catastrophically already for weak mean fields in the limit of large magnetic Reynolds number (i.e.\\ low non-dimensionalized resistivities). Such behavior would imply that mean-field dynamos driven by the $\\alpha$ effect could not generate the observed large-scale magnetic fields. This claim stymied the field of large-scale dynamos for the 1990s. While strong fields are observed in nature, the theoretical understanding appeared to have been cut down. Eventually it was recognized that this behavior is not generally applicable, being restricted to two-dimensional systems, or to homogeneous (non-dynamo generated) mean fields \\citep{BB02}, and large-scale dynamo simulations became common \\citep{B01,BD01}. These new simulations occurred alongside the realization that magnetic helicity conservation, through the dynamical $\\alpha$ quenching formalism, provides an excellent theoretical understanding of the saturation of $\\alpha$--effect dynamos:  the build-up of small-scale magnetic helicity quenches the $\\alpha$ effect \\citep{FB02}. Even so, the question of catastrophic quenching has remained open, with indications of saturated large-scale field strength decreasing with increasing magnetic Reynolds number for shearing sheets and open \\alp systems \\citep{BS05b}.  Further, while the saturation field strength in \\alp systems with periodic or perfectly conducting boundaries has been found to be independent of the resistivity for adequately (and in practice modestly) super-critical $\\Rm$, the timescale to reach saturation increases linearly with $\\Rm$ \\citep{B01}. This has led to the study of magnetic helicity fluxes \\citep{Vishniac, BS04, Mitra10, AdvGauge}, where the hope is that, because the build-up of small-scale magnetic helicity quenches the $\\alpha$ effect, stronger and faster growing dynamos should be possible if the helicity is, instead, exported (as it cannot be destroyed except through the action of true, i.e.\\ microphysical, dissipation).  Probing the reality of catastrophic quenching is naturally difficult.  Analytical theory is impossible, and direct numerical simulations are limited to $\\Rm$ that, while significantly super-critical for many systems, are nevertheless orders of magnitude below those of astrophysical systems. The dynamical $\\alpha$ quenching formalism allows probing large $\\Rm$ in systems it can handle, but its validity there cannot, of course, be directly verified. While the evidence for and against catastrophic quenching is limited, resolving the issue is a crucial step in advancing dynamo theory.  The continued improvement in techniques to measure turbulent dynamo coefficients from simulations has enabled new approaches to evaluating different formulations of the dynamical quenching formalism. In particular, the test-field method \\citep{Sch05,Sch07} has been used to rule out the possibility of catastrophic quenching of the turbulent magnetic diffusivity $\\etat$ in $\\alpha^2$ dynamos \\citep{BRRS08}. Recent advances in the theory of magnetic helicity fluxes in the presence of shear \\citep{Shear} have led us to continue these developments in dynamical quenching by revisiting earlier results from shearing systems. Somewhat surprisingly, these developments return the 1-D $\\alpha$ dependent models to the first 0-D $\\alpha$ dependent models \\citep{BB02}. In addition, we shall extend here earlier numerical studies of \\alp dynamos in open systems. ",
        "conclusions": "We have used the test-field method to examine the predictions of catastrophic $\\alpha$--quenching resulting from dynamically--quenched mean-field models in shearing systems. Formulations for dynamical $\\alpha$--quenching which are superior for the problem of shearing systems do not predict Type 1 catastrophic quenching (reduced field strength) but do predict Type 2 quenching (long final saturation times), extending results that do not allow for spatial variations of $\\alpha$ \\citep{BB02} to models that do. We have further revisited simulations of \\alp dynamos in open systems and, at admittedly quite modest $\\Rm$, found no evidence of field strength scaling inversely with $\\Rm$.  The picture we see now for $\\alpha$--effect dynamos, motivated by the concepts and formalism of dynamical $\\alpha$--quenching, is one of exponential growth during a rapid initial saturation phase. This phase ends when the magnetic helicity in the small-scale fields is comparable to the helicity in the forcing that generates the $\\alpha$--effect. At this point, the total magnetic helicity in the system has not changed from its initial value.  If the system is open, exchanges with the exterior (\\Sec{open}) will tend to keep the \\emph{total} magnetic helicity roughly constant, and the system will then not evolve resistively.  On the other hand, if the system is closed the preferential resistive destruction of magnetic helicity of the small-scale field allows a further resistive growth phase.  It is important to note that the energy in the large-scale field is bounded below by its helicity. Weakly helical large-scale fields are possible, which can have super-equipartition fields even at the end of the kinematic growth phase. Weakly helical large-scale fields are a natural product of sheared system, so rapid growth to sub-equi-, equi- and super-equipartition fields are all expected to occur in nature, although all equi- and super-equipartition fields in the high $\\Rm$ systems of astrophysics are expected to be weakly helical."
    },
    "1107/1107.2431_arXiv.txt": {
        "abstract": "Galaxy star formation rates (SFRs) are sensitive to the local environment; for example, the high-density regions at the cores of dense clusters are known to suppress star formation.  It has been suggested that galaxy transformation occurs largely in groups, which are the intermediate step in density between field and cluster environments. In this paper, we use deep MIPS 24\\micron~observations of intermediate-redshift (0.3 $\\la$ z $\\la$ 0.55) group and field galaxies from the Group Environment and Evolution Collaboration (GEEC) subset of the Second Canadian Network for Observational Cosmology (CNOC2) survey to probe the moderate-density environment of groups, wherein the majority of galaxies are found.  The completeness limit of our study is log(L$_{TIR}$(L$_{\\odot}$)) $\\ga$ 10.5, corresponding to SFR $\\ga$ 2.7 M$_{\\odot}$ yr$^{-1}$.  We find that the group and field galaxies have different distributions of morphologies and mass. However, individual group galaxies have star-forming properties comparable to those of field galaxies of similar mass and morphology; that is, the group environment does not appear to modify the properties of these galaxies directly.  There is a relatively large number of massive early-type group spirals, along with E/S0 galaxies, that are forming stars above our detection limit. These galaxies account for the nearly comparable level of star-forming activity in groups as compared with the field, despite the differences in mass and morphology distributions between the two environments. The distribution of specific SFRs (SFR/M$_{*}$) is shifted to lower values in the groups, reflecting the fact that groups contain a higher proportion of massive and less active galaxies. Considering the distributions of morphology, mass, and SFR, the group members appear to lie between field and cluster galaxies in overall properties. ",
        "introduction": "In a $\\Lambda$CDM universe, galaxies, on average, move from areas of lower density to areas of higher density, merging and combining to form larger and larger systems like the massive galaxy clusters we see in the local universe. Galaxy evolution, therefore, cannot be understood without considering the influence of environment on galaxy properties. The importance of environment on galaxy evolution is demonstrated by the differences between galaxies in clusters and those in the field. For example, the fraction of blue galaxies in clusters has been decreasing since z $\\sim$ 1, and local clusters are dominated by red galaxies (Butcher \\& Oemler 1978; Kennicutt 1983; Hashimoto et al. 1998; Andreon et al. 2004; Poggianti et al. 2006; Cooper et al. 2007; Loh et al. 2008; Cucciati et al. 2010). Dressler (1980) found a dramatic increase in the proportion of early-type galaxies with local density inside rich clusters, i.e., the morphology-density relation. This relation has also been shown to extend down to the group environment (Postman \\& Geller 1984). Most of these cluster early-type galaxies have not had appreciable star formation in gigayears, though as we move outward from the centers of the clusters, we see more late-type galaxies overall and more early-type galaxies that have had more recent star formation (Balogh et al. 1997, 1999; Bai et al. 2009).  Environment is expected to influence the rates of both gas exhaustion and interactions.  The drop in star-forming activity in dense environments might be due to interactions with the inter-galactic gas, such as through ram-pressure stripping of cold gas in galaxies (Gunn \\& Gott 1972; Larson et al. 1980; Kinney et al. 2004) or stripping of the hot gas through strangulation (Balogh et al. 2000; Kawata \\& Mulchaey 2008; McCarthy et al. 2008). Alternatively, the lower fraction of star forming galaxies might arise through galaxy-galaxy interactions, either major mergers or harassment (frequent high-speed encounters of galaxies that do not lead to mergers), both of which would accelerate the star formation and lead to early exhaustion of the interstellar material. Tides raised by the overall gravitational potential of dense clusters may also play a role (Henriksen \\& Byrd 1996). It has been proposed that the morphological transformation into early-type galaxies may occur first, suppressing star formation by stabilizing the gas against fragmentation (Martig et al. 2009), although other recent studies question this claim \\citep{kovac10}.  Explaining the behavior of galaxies in different environments in terms of consistent theories for the growth of galaxies in the early universe has proven challenging (Bower et al. 2006; Kaviraj et al. 2009). Therefore, an intense observational approach is needed to help develop an understanding of the relation between environment and galaxy evolution. It should be possible to disentangle environmental influences by comparing galaxy behavior in different environments, though just as the mechanisms for galaxy transformation are not fully understood, the environments where these processes operate also need to be explored and defined. Additionally, most of the previous work regarding star formation with respect to environment focused on rich clusters (e.g., Dressler et al. 2009; Bai et al. 2009; Haines et al. 2009, and references therein); relatively few studies have focused on groups (e.g. Zabludoff \\& Mulchaey 1998, 2000; Balogh et al. 2004; Johnson et al. 2007; Marcillac et al. 2008; Bai et al. 2010).  Simard et al. (2009) argue that cluster-centric processes are not the dominant factor in galaxy morphological transformation.  The majority of galaxies live in the less-dense group environment (Geller \\& Huchra 1983; Eke et al 2004), and because clusters probably form from coalescing groups and field galaxies, much of the evolution apparent in cluster galaxies may have occurred in groups prior to their assimilation into clusters. Zabludoff \\& Mulchaey (1998) found that the proportion of early-type galaxies in groups ranged from that typical of the field ($\\sim$25\\%) to that found in dense clusters ($\\sim$55\\%), suggesting that much of the morphological transformation of galaxies from field to cluster properties occurs in groups. \\citet{just10} show that an increase in the proportion of S0 galaxies with decreasing redshift occurs in moderate-mass groups/poor clusters ($\\sigma < 750$ km s$^{-1}$), and \\citet{wilman09} find that for groups at intermediate redshifts, the fraction of S0s is as high as in clusters, even at fixed luminosities. Additionally, because most galaxies reside in groups---and, therefore, galaxies spend most of their time in groups---uncovering the effects of these lower-density environments can help us understand the evolutionary path of the global galaxy population over cosmic time \\citep{mcgee09}.  If morphological transformations frequently occur at intermediate densities, is this also the location where star formation is cutoff? Previous studies of star formation in groups have mostly focused on optical indicators such as H$\\alpha$ or [OII]$\\lambda$3727 emission lines or ultraviolet continuum. These measures can extend to low levels of star formation, and they indicate a significantly lower level of activity in groups and clusters than in the field (Wilman et al. 2005a; Gerke et al. 2007; Balogh et al. 2009; Iovino et al. 2010; Peng et al. 2010). Such first-order comparisons are usually made under the assumption that the extinction is similar in groups and clusters (and that star formation is not deeply obscured). However, corrections are required to convert these optical indicators to accurate star formation rates (SFRs).  At z $\\ga$ 0.3, H$\\alpha$ moves out of the range of optical spectroscopy and SFR estimates rely on [OII], where the extinction is large and uncertain. The [OII]$\\lambda3727$ line has additional problems of being sensitive to dust reddening and metallicity, thus being relatively weak in high-mass systems, although empirical corrections have been suggested to correct for this and other effects (Moustakas et al. 2006; Gilbank et al. 2010).  The infrared (IR) is advantageous for probing high levels of star formation. While SFRs determined in the IR are only sensitive to dust-obscured star formation, the correction to the total star formation in objects at moderate to high luminosities (L$_{TIR} \\ga 1 \\times 10^{10}$ L$_{\\odot}$) is small \\citep{rieke09}. Nonetheless, previous IR studies have not reached consistent conclusions about star formation in groups.  \\citet{marcillac08} found no significant dependence on the incidence of luminous infrared galaxies (LIRGs) as a function of field or group environment at z $\\sim$ 0.8.  At lower redshifts, \\citet{wilman08} find a dearth of star formation in group galaxies at z $\\sim$ 0.4, while \\citet{tran09} find a similar incidence of SFRs in massive groups (for galaxies forming stars $> 3 M_\\odot$ yr$^{-1}$) as in the field at z $\\sim$ 0.37.  For local groups, Bai et al. (2010) report rates of star formation somewhat lower than in the field (by $\\sim$30\\%).  This paper is a step toward understanding the influence of the group environment on star formation and resolving some of the apparent discrepancies in previous studies of the same topic.  Some of these discrepancies, especially in cluster studies, may arise from ``field'' samples contaminated by groups.  To compare groups with actual field galaxies, we need a clean field sample with as many galaxies as possible. Here, we present 24\\micron~measurements of 232 group galaxies and 236 field galaxies from 0.3 $\\la$ z $\\la$ 0.55 in the Second Canadian Network for Observational Cosmology (CNOC2) survey.  The paper is organized as follows. In Section 2, we discuss the sample, data reductions, and errors.  Section 3 covers the construction of fractional IR luminosity functions (LFs) and compares the IR luminosities, morphologies, and masses of group and field galaxies to estimate what effect (if any) environment has on star formation at these redshifts.  We discuss the implications of our results in Section 4. For all cosmological corrections, we assume the parameters $H_0 = $70~km~s$^{-1}$Mpc$^{-1}$,~$\\Omega_M = 0.3,~\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "We have observed 26 galaxy groups and accompanying field galaxies with deep MIPS photometry and used 24\\micron~flux densities to estimate the total IR luminosities and SFRs of galaxies in both environments.  We find that on an individual basis, group and field galaxies of similar mass and morphology do not differ significantly in terms of their SFRs, and the amount of star formation does not depend on the richness of the groups.  However, the groups have systematically lower sSFRs and a higher incidence of massive early-type galaxies, more reminiscent of clusters than the field.  We discovered that some of these E and S0 galaxies, as well as a large contingent of massive early spirals, are still forming stars at significant levels.  These galaxies may explain why the fLFs of the groups and field are nearly identical despite the overall decrease in star-forming activity in the groups. The group environment affects galaxy SFRs primarily through the shift toward higher masses, with an accompanying trend toward earlier types and reduced sSFRs. These high-mass, early-type galaxies, along with IR luminosities comparable to the field, put groups in between the field and clusters in terms of overall galaxy properties."
    },
    "1107/1107.0744_arXiv.txt": {
        "abstract": "We report the discovery of a threshold in the H{\\sc i} column density of Galactic gas clouds below which the formation of the cold phase of H{\\sc i} is inhibited. This threshold is at ${\\rm N_{\\sc HI}} = 2 \\times 10^{20}$~per~cm$^{2}$; sightlines with lower \\hi\\ column densities have high spin temperatures (median ${\\rm T}_s \\sim 1800$~K), indicating low fractions of the cold neutral medium (CNM), while sightlines with ${\\rm N_{\\sc HI}} \\ge 2 \\times 10^{20}$~per~cm$^{2}$ have low spin temperatures (median ${\\rm T}_s \\sim 240$~K), implying high CNM fractions. The threshold for CNM formation is likely to arise due to inefficient self-shielding against ultraviolet photons at lower \\hi\\ column densities. The threshold is similar to the defining column density of a damped Lyman-$\\alpha$ absorber; this indicates a physical difference between damped and sub-damped Lyman-$\\alpha$ systems, with the latter class of absorbers containing predominantly warm gas. ",
        "introduction": "\\label{sec:intro}  The diffuse interstellar medium (ISM) contains gas over a wide range of densities, temperatures and ionization states. These can be broadly sub-divided into the molecular phase (e.g. \\citealt{snow06}), the neutral atomic phase (e.g. \\citealt{kulkarni88}), and the warm and hot ionized phases (e.g. \\citealt{haffner09}). The neutral atomic medium (mostly neutral hydrogen, \\hi) is further usually sub-divided into ``cold'' and ``warm'' phases (the ``CNM'' and ``WNM'', respectively). Typical CNM temperatures and densities are $\\sim 40-200$~K and $\\gtrsim 10$~cm$^{-3}$, respectively, with corresponding WNM values of $\\gtrsim 5000$~K and $\\sim 0.1-1$~cm$^{-3}$. This was originally an observational definition, to distinguish between phases producing strong narrow absorption lines towards background radio-loud quasars and smooth broad emission lines \\citep{clark65}. Later, this separation into cold and warm phases was found to arise naturally in the context of models in which the atomic and ionized phases are in pressure equilibrium.  Atomic gas at intermediate temperatures was found to be unstable to perturbations, and to migrate to either the cold or warm stable phases, given sufficient time to attain pressure equilibrium \\citep{field69,wolfire95}.  Radio studies of the \\hi-21cm transition have played a vital role in understanding physical conditions in the neutral ISM. The \\hi-21cm excitation temperature of a gas cloud (the ``spin temperature'', $\\ts$) can be shown to depend on the local kinetic temperature ($\\tk$; \\citealt{field59,liszt01}), with $\\ts \\approx \\tk$ in the CNM and $\\ts \\le \\tk$ in the WNM \\citep{liszt01}. Since the \\hi-21cm absorption opacity for a fixed \\hi\\ column density varies inversely with $\\ts$, while the emissivity is independent of $\\ts$ (in the low opacity limit), \\hi-21cm absorption studies against background radio sources are primarily sensitive to the presence of CNM along the sightline, while \\hi-21cm emission studies are sensitive to both warm and cold \\hi. Indeed, the original evidence for two temperature phases in the neutral gas stemmed from a comparison between \\hi-21cm absorption and emission spectra \\citep{clark65}.  Physical conditions in the neutral gas are also known to depend on the gas column density. At the low \\hi\\ column densities of the inter-galactic medium, $\\nhi < 10^{17}$~\\cm, the gas is optically thin to ionizing ultraviolet (UV) photons and is hence mostly ionized, giving rise to the so-called Lyman-$\\alpha$ forest \\citep{rauch98}. At higher \\hi\\ column densities, $10^{17} \\lesssim \\nhi \\lesssim 10^{20}$~\\cm, the \\hi\\ is optically-thick to ionizing photons at the Lyman-limit, and the core of the Lyman-$\\alpha$ absorption line is saturated. Such ``Lyman-limit systems'' are partially ionized and typically arise for sightlines through the outskirts of galaxies \\citep{bergeron91}.  At still higher \\hi\\ column densities, $\\nhi \\ge 2 \\times 10^{20}$~\\cm, the Lyman-$\\alpha$ line becomes optically-thick in its naturally-broadened wings and acquires a Lorentzian shape; such systems are called damped Lyman-$\\alpha$ absorbers (DLAs; \\citealt{wolfe05}). Damped absorption is ubiquitous on sightlines through galaxy disks and, at high redshifts, has long been used as the signature of the presence of an intervening galaxy along a quasar sightline \\citep{wolfe86}. Finally, the molecular hydrogen fraction shows a steep transition from very low values to greater than 5\\% at $\\nhi \\sim 5 \\times 10^{20}$~\\cm\\ in the Galaxy \\citep{savage77,gillmon06}). Neutral gas is likely to become predominantly molecular at higher \\hi\\ column densities, $\\nhi \\gtrsim 10^{22}$~\\cm\\ \\citep{schaye01,krumholz09b}.  While it is well known that a threshold column density ($\\nhi \\sim 5 \\times 10^{20}$~\\cm) is required to form the molecular phase, for self-shielding against UV photons in the ${\\rm H}_2$ Lyman band \\citep{stecher67,hollenbach71,federman79}, the conditions for CNM formation are less clear. In this {\\it Letter}, we report results from \\hi-21cm absorption studies of a large sample of compact radio sources that indicate the presence of a similar, albeit lower, \\hi\\ column density threshold for CNM formation in the Milky Way. ",
        "conclusions": "} \\begin{tabular}{|c|c|c|c|c|} \\hline QSO        & Latitude & N$_{\\rm HI}\\:^a$         &   $\\int \\tau \\rm{dV}\\:^b$  &     $\\ts$        \\\\ &  $b$     & $\\times 10^{20}$ \\cm &        \\kms\\           &       K          \\\\ \\hline B0023-263  & -84.17  &  $  1.641  \\pm  0.022   $ &  $ 0.025  \\pm  0.005  $ &  $  3546  \\pm  694  $  \\\\ B0114-211  & -81.47  &  $  1.380  \\pm  0.034   $ &  $ 0.135  \\pm  0.005  $ &  $   559  \\pm   26  $  \\\\ B0117-156  & -76.42  &  $  1.418  \\pm  0.034   $ &  $ 0.031  \\pm  0.004  $ &  $  2538  \\pm  353  $  \\\\ B0134+329  & -28.72  &  $  4.272  \\pm  0.029   $ &  $ 0.443  \\pm  0.002  $ &  $   530  \\pm    4  $  \\\\ B0202+149  & -44.04  &  $  4.809  \\pm  0.027   $ &  $ 0.747  \\pm  0.005  $ &  $   353  \\pm    3  $  \\\\ B0237-233  & -65.13  &  $  2.078  \\pm  0.038   $ &  $ 0.294  \\pm  0.004  $ &  $   388  \\pm    9  $  \\\\ B0316+162  & -33.60  &  $  9.431  \\pm  0.034   $ &  $ 2.964  \\pm  0.004  $ &  $   175  \\pm    1  $  \\\\ B0316+413  & -13.26  &  $ 13.234  \\pm  0.028   $ &  $ 1.941  \\pm  0.003  $ &  $   374  \\pm    1  $  \\\\ B0355+508  &  -1.60  &  $ 74.344  \\pm  0.079   $ &  $ 45.820  \\pm  1.120 $ &  $    89  \\pm    2  $  \\\\ B0404+768  &  18.33  &  $ 10.879  \\pm  0.032   $ &  $ 1.945  \\pm  0.005  $ &  $   307  \\pm    1  $  \\\\ B0407-658  & -40.88  &  $  3.363  \\pm  0.018   $ &  $ 0.548  \\pm  0.007  $ &  $   337  \\pm    5  $  \\\\ B0429+415  &  -4.34  &  $ 37.067  \\pm  0.041   $ &  $ 10.879  \\pm  0.007 $ &  $   187  \\pm    1  $  \\\\ B0438-436  & -41.56  &  $  1.380  \\pm  0.027   $ &  $ < 0.020            $ &  $  > 3785 $  \\\\ B0518+165  & -11.34  &  $ 20.616  \\pm  0.035   $ &  $ 6.241  \\pm  0.007  $ &  $   181  \\pm    1  $  \\\\ B0531+194  &  -7.11  &  $ 26.609  \\pm  0.036   $ &  $ 4.062  \\pm  0.005  $ &  $   359  \\pm    1  $  \\\\ B0538+498  &  10.30  &  $ 19.542  \\pm  0.029   $ &  $ 5.618  \\pm  0.003  $ &  $   191  \\pm    1  $  \\\\ B0831+557  &  36.56  &  $  4.469  \\pm  0.032   $ &  $ 0.483  \\pm  0.006  $ &  $   507  \\pm    7  $  \\\\ B0834-196  &  12.57  &  $  7.125  \\pm  0.036   $ &  $ 0.973  \\pm  0.005  $ &  $   402  \\pm    3  $  \\\\ B0906+430  &  42.84  &  $  1.251  \\pm  0.036   $ &  $ 0.051  \\pm  0.003  $ &  $  1342  \\pm   89  $  \\\\ B1151-348  &  26.34  &  $  7.732  \\pm  0.031   $ &  $ 0.714  \\pm  0.005  $ &  $   594  \\pm    5  $  \\\\ B1245-197  &  42.88  &  $  3.802  \\pm  0.037   $ &  $ 0.158  \\pm  0.005  $ &  $  1323  \\pm   40  $  \\\\ B1323+321  &  81.05  &  $  1.260  \\pm  0.034   $ &  $ 0.083  \\pm  0.003  $ &  $   836  \\pm   40  $  \\\\ B1328+254  &  80.99  &  $  1.065  \\pm  0.036   $ &  $ 0.021  \\pm  0.003  $ &  $  2729  \\pm  411  $  \\\\ B1328+307  &  80.67  &  $  1.197  \\pm  0.030   $ &  $ 0.072  \\pm  0.002  $ &  $   916  \\pm   32  $  \\\\ B1345+125  &  70.17  &  $  1.957  \\pm  0.024   $ &  $ 0.305  \\pm  0.005  $ &  $   352  \\pm    7  $  \\\\ B1611+343  &  46.38  &  $  1.318  \\pm  0.034   $ &  $ 0.019  \\pm  0.003  $ &  $  3873  \\pm  565  $  \\\\ B1641+399  &  40.95  &  $  1.044  \\pm  0.032   $ &  $ 0.009  \\pm  0.002  $ &  $  6498  \\pm 1760  $  \\\\ B1814-637  & -20.76  &  $  6.416  \\pm  0.018   $ &  $ 0.997  \\pm  0.007  $ &  $   353  \\pm    3  $  \\\\ B1827-360  & -29.34  &  $  8.163  \\pm  0.018   $ &  $ 1.542  \\pm  0.003  $ &  $   290  \\pm    1  $  \\\\ B1921-293  & -11.78  &  $  7.312  \\pm  0.017   $ &  $ 1.446  \\pm  0.006  $ &  $   277  \\pm    1  $  \\\\ B2050+364  & -48.84  &  $ 27.638  \\pm  0.038   $ &  $ 3.024  \\pm  0.010  $ &  $   501  \\pm    2  $  \\\\ B2200+420  &  -5.12  &  $ 17.119  \\pm  0.042   $ &  $ 3.567  \\pm  0.016  $ &  $   263  \\pm    1  $  \\\\ B2203-188  & -51.16  &  $  2.432  \\pm  0.018   $ &  $ 0.248  \\pm  0.004  $ &  $   537  \\pm   10  $  \\\\ B2223-052  & -10.44  &  $  4.561  \\pm  0.030   $ &  $ 1.034  \\pm  0.003  $ &  $   242  \\pm    2  $  \\\\ B2348+643  &   2.56  &  $ 70.522  \\pm  0.051   $ &  $ 32.517  \\pm  0.031 $ &  $   119  \\pm    1  $  \\\\ &         &                        &                        &                   \\\\ \\hline \\end{tabular} \\vskip 0.01in Notes: $^a$~From the LAB emission survey. $^b$~From our \\hi-21cm absorption survey. \\end{center} \\end{table}   \\label{sec:discuss}  \\begin{figure}[t!] \\epsfig{file=fig2.eps,height=3.5truein,width=3.5truein} \\caption{The column-density-weighted harmonic mean spin temperature (K) plotted against Galactic latitude. } \\label{fig:fig2} \\end{figure}  Most theoretical models of the diffuse interstellar medium treat it as a multi-phase medium, consisting of neutral and ionized phases in pressure equilibrium (e.g. \\citealt{field69,mckee77,wolfire95}). In the McKee-Ostriker model, physical conditions in the ISM are regulated by supernova explosions, whose blast waves sweep up gas in the ISM into shells, leaving large cavities full of hot ionized gas \\citep{mckee77}. Cold atomic gas is then formed by the rapid cooling of the swept-up and shocked gas. Soft X-rays from neighbouring hot ionized gas and stellar UV photons penetrate into the CNM and heat and partially ionize it, producing envelopes of warm neutral and warm ionized gas (WIM). Thus, a typical ISM ``cloud'' [see Fig.~1 of \\citet{mckee77}] is expected to form at the edges of supernova remnants and super-bubbles containing the hot ionized medium (HIM). Such a cloud consists of a CNM core surrounded by a WNM envelope and a further WIM shell, all of which are in pressure equilibrium. The CNM core is almost entirely neutral, the WNM and WIM are partially ionized, and the HIM is almost entirely ionized \\citep{mckee77}.  A threshold \\hi\\ column density for CNM formation in an ISM cloud arises quite naturally due to the need for self-shielding against ionizing ultraviolet (and soft X-ray) photons (see also \\citealt{schaye04}). At low \\hi\\ columns, UV photons can penetrate all the way into a cloud core and heat (and partly ionize) the \\hi. Only clouds that self-shield against the penetration of these high-energy photons can retain a stable CNM core. Self-shielding against UV photons is significant once the optical depth at the Lyman limit crosses unity, i.e. for $\\nhi \\ge 10^{17}$~\\cm, and becomes more and more efficient with increasing \\hi\\ column density. Our results indicate that self-shielding only becomes entirely efficient at excluding UV photons from the interior of \\hi\\ clouds at a {\\it total} \\hi\\ column density of $\\nhi \\sim 2 \\times 10^{20}$~\\cm.  Below this threshold, sufficient UV photons penetrate into the cloud interior to hinder the survival of the cold phase. The simple ``sandwich'' model overlaid on Fig.~\\ref{fig:fig1}[A] suggests that a shielding column of $\\sim (0.5 - 0.75) \\times 10^{20}$~\\cm\\ {\\it on each exposed surface of a cloud} is sufficient to exclude UV photons, with the rest of the \\hi\\ then free to cool to lower temperatures.  It is also possible that sightlines with low $\\nhi$ are sampling gas at higher distances from the Galactic plane, with lower metallicity and pressure.  Two-phase models that incorporate detailed balancing of heating and cooling rates obtain lower CNM fractions at lower pressures and metallicities \\citep{wolfire95}.  Observationally, \\citet{kanekar09c} have found evidence for an anti-correlation between spin temperature and metallicity in high-$z$ DLAs, with low-metallicity DLAs having higher spin temperatures, probably because the paucity of metals yields fewer cooling routes \\citep{kanekar01a}.  If low-$\\nhi$ sightlines contain ``clouds'' with systematically lower metallicities (and/or pressures), these could have lower cooling rates and hence, lower CNM fractions. Unfortunately, we do not have estimates of either pressure or metallicity along our sightlines and hence cannot test this possibility. We note that all sightlines with high spin temperatures are at intermediate or high Galactic latitudes, $b > 40^\\circ$ (see Fig.~\\ref{fig:fig2}), suggesting larger distances from the plane.  Conversely, the figure also shows that low spin temperatures ($\\lesssim 400$~K) are obtained even for sightlines at high Galactic latitudes ($b \\sim 40-70^\\circ$). Measurements of the metallicity and pressure along the low-$\\nhi$ sightlines, through UV spectroscopy, would be of much interest.  There have been earlier suggestions, based on semi-analytical or numerical treatments of self-shielding, that \\hi\\ clouds are predominantly neutral for $\\nhi \\gtrsim 10^{20}$~\\cm\\ (e.g. \\citealt{viegas95,wolfe05}). However, this is the first direct evidence for a change in physical conditions in \\hi\\ clouds at this column density. Note that this is a factor of several lower than the known threshold of $\\nhi \\sim 5 \\times 10^{20}$~\\cm\\ for the formation of molecular hydrogen in Galactic clouds, set by the requirement of self-shielding against UV photons at wavelengths in the H$_2$ Lyman band \\citep{savage77}. Thus, there appear to be three column densities at which phase transitions occur in ISM clouds, at $\\nhi \\sim 2 \\times 10^{20}$~\\cm\\ resulting in the formation of cold \\hi, at $\\nhi \\sim 5 \\times 10^{20}$~\\cm\\ resulting in the formation of molecular hydrogen, and finally, at $\\nhi > 10^{22}$~\\cm, when most of the atomic gas is converted into the molecular phase.  In this context, it is vital to appreciate that the abscissa of Fig.~\\ref{fig:fig1} refers to {\\it apparent\\ } \\hi\\ column density under the assumption of negligible self-opacity of the \\hi\\ profile. The saturation column density, N$_\\infty$, is an {\\it observational} upper limit to $\\int T_B dV$ and does not represent a physical limit on $\\nhi$. In fact, the widespread occurrence of the \\hi\\ self-absorption phenomenon within the Galaxy \\citep{gibson05} and the detailed modeling of high resolution extragalactic \\hi\\ spectra suggests that self-opacity is a common occurrence \\citep{braun09} which can disguise neutral column densities that reach $\\nhi \\sim 10^{23}$~\\cm.  The original definition of a DLA as an absorber with $\\nhi \\ge 2 \\times 10^{20}$~\\cm\\ was an observational one, based on the requirement that the damping wings of the Lyman-$\\alpha$ line be detectable in low-resolution optical spectra of moderate sensitivity \\citep{wolfe86}.  With today's 10-m class optical telescopes, it is easy to detect the damping wings at significantly lower \\hi\\ column densities, $\\sim 10^{19}$~\\cm (e.g. \\citealt{peroux03}); such systems are referred to as ``sub-DLAs''. There has been much debate in the literature on whether or not DLAs and sub-DLAs should be treated as a single class of absorber and on their relative importance in contributing to the cosmic budget of neutral hydrogen and metals (e.g. \\citealt{peroux03,wolfe05,kulkarni10}).  \\citet{wolfe05} argue that DLAs and sub-DLAs are physically different, claiming that most of the \\hi\\ in sub-DLAs is ionized and at high temperature, while that in DLAs is mostly neutral; this is based on numerical estimates of self-shielding in DLAs and sub-DLAs against the high-$z$ UV background \\citep{viegas95}. Our detection of an \\hi\\ column density threshold for CNM formation that matches the defining DLA column density indicates that DLAs and sub-DLAs are indeed physically different types of absorbers, with sub-DLAs likely to have significantly lower CNM fractions than DLAs, at a given metallicity.  In summary, we report the discovery of a threshold \\hi\\ column density, $\\nhi \\sim 2 \\times 10^{20}$~\\cm, for cold gas formation in \\hi\\ clouds in the interstellar medium. Above this threshold, the majority of Galactic sightlines have low spin temperatures, $\\ts \\lesssim 500$~K, with a median value of $\\sim 240$~K. Below this threshold, typical sightlines have far higher spin temperatures, $> 600$~K, with a median value of $\\sim 1800$~K. The threshold for CNM formation appears to arise naturally due to the need for self-shielding against ultraviolet photons, which penetrate into the cloud at lower \\hi\\ columns and heat and ionize the \\hi, inhibiting the formation of the cold neutral medium."
    },
    "1107/1107.3437_arXiv.txt": {
        "abstract": "In the last decades, force-free-field modelling has been used extensively to describe the coronal magnetic field and to better understand the physics of solar eruptions at different scales. Especially the evolution of active regions has been studied by successive equilibria in which each computed magnetic configuration is subject to an evolving photospheric distribution of magnetic field and/or electric current density. This technique of successive equilibria has been successful in describing the rate of change of the energetics for observed active regions. Nevertheless the change in the magnetic configuration due to the increase/decrease of electric current for different force-free models (potential, linear and nonlinear force-free fields) has never been studied in detail before. Here we focus especially on the evolution of the free magnetic energy, the location of the excess of energy, and the distribution of electric currents in the corona. For this purpose, we use an idealised active region characterised by four main polarities and a satellite polarity allowing us to specify a complex topology and sheared arcades to the coronal magnetic field but no twisted flux bundles. We investigate the changes in the geometry and connectivity of field lines, the magnetic energy and current density content as well as the evolution of null points. Increasing the photospheric current density in the magnetic configuration does not dramatically change the energy-storage processes within the active region even if the magnetic topology is slightly modified. We conclude that for reasonable values of the photospheric current density (the force-free parameter $\\alpha < 0.25$ Mm$^{-1}$), the magnetic configurations studied do change but not dramatically: {\\em i}) the original null point stays nearly at the same location, {\\em ii}) the field-line geometry and connectivity are slightly modified, {\\em iii}) even if the free magnetic energy is significantly increased, the energy storage happens at the same location. This extensive study of different force-free models for a simple magnetic configuration shows that some topological elements of an observed active region, such as null points, can be reproduced with confidence only by considering the potential-field approximation. This study is a preliminary work aiming at understanding the effects of electric currents generated by characteristic photospheric motions on the structure and evolution of the coronal magnetic field. ",
        "introduction": "\\label{sec:intro}  One key unsolved issue in solar physics is the generation and effects of electric currents from photospheric motions of the stressed and sheared coronal magnetic field. With the development of reliable techniques such as coronal magnetic-field extrapolations based on complex distributions of photospheric currents  (see {\\em e.g.}, reviews by \\citeauthor{reg07b}, \\citeyear{reg07b}; \\citeauthor{wie08}, \\citeyear{wie08}), it is important to understand how a modelled magnetic-field configuration is subject to change due to slight modifications of the photospheric-current distribution and thus the difference between several force-free assumptions using the same boundary conditions. The main aim is to understand these changes in the geometry and topology of field lines assuming that the magnetic field is in a force-free equilibrium. In \\inlinecite{reg07}, we have performed a first comparison between four different active regions with different behaviours due to the complex distribution of the photospheric field and of the electric currents representing the history of the evolution of the active region. This comparison was done for the same force-free model, namely the nonlinear force-free field. We showed that, statistically speaking, the magnetic-field lines are longer and higher in a nonlinear force-free field compared to the corresponding potential field. To develop our understanding of the effects of electric currents on magnetic configurations, \\inlinecite{reg09a} compared the behaviour of a simple bipolar field subject to different distributions of electric currents using different force-free models. The bipolar field has been studied in terms of magnetic-energy storage and magnetic-helicity changes. We showed that the amount of electric currents that can be injected in a magnetic-equilibrium configuration depends strongly on the spatial distribution of the currents, the existence of return currents having a stabilising effect on the magnetic configuration. The next step developed in this article is to understand the effects of electric currents on a magnetic configuration having predominant topological elements ({\\em i.e.} a null point in the domain of interest).  In the past decades, magnetic topology has become a key ingredient in understanding the origin of flares in active regions.  The general definition of magnetic topology concerns the properties of magnetic-field lines and magnetic-flux surfaces that are invariant under continuous deformation in plasma conditions satisfying the frozen-in assumption \\cite{low06, low07, jan10a, ber06}. This definition implies that the number of null points does not change and the connectivity of field lines rigidly anchored to the boundaries of the domain is also invariant under the above conditions. However, we focus here on the evolution of topological elements and connectivity of field lines obtained from force-free models: the field lines are not rigidly anchored (the models allow for reconnection of field lines) in order to maintain a force-free state. Since the magnetic topology is not required to be preserved, we restrict the definition of the magnetic topology for a given equilibrium state to the ensemble of topological elements forming the magnetic skeleton \\cite{bun96}. Despite an extensive literature on magnetic topology, the theoretical background in three dimensions was developed only recently (see review by \\opencite{pri00}). The 3D topological elements constituting the skeleton of a magnetic configuration can be divided into two parts: the true topology containing null points, separatrix surfaces, and separators, and the quasi-topology including quasi-separatrix layers and hyperbolic flux tubes (and the true topology). We focus especially on the location and properties of null points as a proxy for describing the topology of magnetic configuration. This study aims at understanding the possible changes of properties and location of null points subject to the continuous variation of a free parameter in various force-free models.  It has been proven that the topological elements of a coronal magnetic configuration are of prime importance to study the onset of flares and coronal mass ejections (CMEs). For instance, in the classical model of flare, magnetic reconnection occurs at a null point or in a current sheet formed along a topological element. Combining observations and coronal-field models, \\citeauthor{aul00} have shown that a powerful flare associated with a CME involves a coronal null point and a spine field line. This topological study in conjunction with EUV observations supports the breakout model \\cite{ant99} as a triggering mechanism for this particular event. Recently, \\inlinecite{zhao08} have derived the skeleton of an active region and its temporal evolution before and after an eruptive event. The authors have found several coronal null points in a quasi-force-free field configuration. Unfortunately, the null points found by \\inlinecite{zhao08} did not satisfy the properties of null points for force-free fields and for divergence-free fields (see Appendix A). Other topological studies have been carried out to better understand the release of magnetic energy and the reconnection processes in active region evolution \\cite{dem94,den05,bar05,reg06,li06,luo07,bar07}, during blinkers  \\cite{sub08}, and in the quiet Sun \\cite{schr02,clo04,reg08a}  Our work has been motivated by two earlier articles on this topic by \\inlinecite{dem94} and \\inlinecite{hud99}. \\inlinecite{dem94} have compared the topology of potential and linear force-free fields for a quadrupolar configuration with a coronal null. They found that the topology is similar for the point-charge model but they also noticed that for a bipolar model (based on extended sources) another null can be created in the linear force-free configuration for large value of the force-free parameter $\\alpha$. Based on the point-charge model, \\inlinecite{bun96} reached the same conclusion for different charge distributions. \\inlinecite{hud99} have considered a quadrupolar point charge, symmetric distribution to study the connectivity of field lines for potential, linear, and nonlinear force-free fields. They concluded that the topology can be insanely different from one model to the other. For non-symmetric cases, \\inlinecite{bro00} have found that the topology of a force-free configuration can be similar using the point-charge model. As a step forward to the understanding of the magnetic topology of reconstructed coronal fields, we carry out a comparison between different models of magnetic fields (potential, linear, and nonlinear force-free fields) for a configuration having a coronal null point assuming a continuous distribution of magnetic field at the bottom boundary and no symmetry. We describe the changes in the magnetic configurations based on the evolution of the geometry and connectivity of magnetic field lines and in terms of distribution of magnetic energy and electric currents. In addition, we study the properties of null points as a proxy of the complexity of the field, keeping in mind that the separators and separatrices play an important role in the release of magnetic energy ({\\em e.g.}, \\citeauthor{pri05}, \\citeyear{pri05}). It is worth noticing that we only focus on the modelling of magnetic-field equilibria and their changes through the increase of electric currents, whilst recently \\inlinecite{san11} have described the changes in magnetic fields subject to characteristic photospheric motions using a magnetohydrodynamic (MHD) approach. They have found that the topology of a quadrupolar magnetic field remains stable whatever the perturbations imposed even in the case of strong currents. The  authors claimed that their results can easily be generalised to more complex magnetic-field distributions and non-generic and symmetric cases. Despite their sophisticated MHD approach, they do not study the changes in magnetic energy or in the connectivity of field lines.  In Section~\\ref{sec:mk_np}, we construct a magnetic configuration with a coronal null point from which we will reconstruct the different models (see Section~\\ref{sec:model}). We thus analyse the geometry of field lines in Section~\\ref{sec:geo} and their connectivity in Section~\\ref{sec:con}. The change in topology in the different models is discussed in Section~\\ref{sec:topo}. In addition, we study how the magnetic energy is stored (Section~\\ref{sec:nrj}) and the electric currents are concentrated (Section~\\ref{sec:cur}). In Section~\\ref{sec:concl}, we discuss the implications for future topological studies from reconstructed magnetic fields.  \\begin{figure} \\centering \\includegraphics[width=.531\\linewidth]{sregnier_fig1a.eps} \\includegraphics[width=.459\\linewidth, bb= 130 100 790 860, clip=]{sregnier_fig1b.eps} \\caption{(Left) Quadrupolar distribution of the vertical component of the magnetic field ($B_z$) used as initial field (black and white are negative (N1 and N2) and positive (P1 and P2) polarities). The total magnetic flux is balanced. (Right) Close-up of a few field lines (red and green) depicting the geometry and topology of the potential-field configuration. White polarities and red contours (black polarities and blue contours) are positive ( negative) values of $B_z$. } \\label{fig:bz_quadru} \\end{figure} ",
        "conclusions": "\\label{sec:concl}  We investigated the changes in the magnetic-field configurations obtained for different force-free models (potential, linear, and nonlinear force-free fields) using the same boundary conditions. We analysed the changes in terms of geometry and connectivity of field lines, magnetic energy, and electric-current distributions. We performed this analysis for a continuous magnetic-field distribution with no symmetry. We imposed an electric-current distribution that we have proven to be stable when a large amount of current is injected into the magnetic configuration \\cite{reg09a}. Despite the previous works on this topic, we have here provided an extended analysis of magnetic configurations with a topology that has never been performed before.  For this experiment, the initial configuration corresponds to a potential field with five sources (two large bipoles and one parasitic polarity) having a null point NP0 in the corona. By injecting currents in the magnetic configuration, the geometry and topology of the linear and nonlinear force-free configurations are modified such that:  \\begin{itemize} \\item{the geometry of the field lines is modified similarly to previous results obtained on the same topic: statistically, the field lines are higher and longer when the absolute value of electric current is increased;} \\item{the connectivity of the field lines can be strongly modified near the topological elements (where the connectivity is changed rapidly);} \\item{the initial null point NP0 is moved slightly up or down when the force-free parameter ($\\alpha$) varies but remains with the same basic properties, in particular the spectral radius remains almost constant;} \\item{other null points (up to five) can appear in the magnetic configurations; most of them are located near the boundaries but one (NP1) which propagates towards the strong-field region when the current density is increased;} \\item{the magnetic energy and current distributions can highlight the location of stable null points where strong magnetic-field gradients are present.} \\end{itemize}  We also noticed that for reasonable values of the electric current injected or the force-free parameter $\\alpha$ ($<$ 0.25 Mm$^{-1}$ in this experiment) the magnetic configuration is almost not modified compared to the potential-field configuration.  We thus state that null points existing in potential-field configurations are also present in force-free configurations with the same properties ({\\em e.g.}, sign, spectral radius). This statement means that the null points found in a potential field with a large spectral radius can be considered as stable null points in other magnetic-field models. Even if true for this experiment we need to confirm this statement in a future statistical study of solar active region magnetic fields. It is important to note here that this statement is true {\\em i}) when the force-free fields are computed with the same boundary conditions, {\\em ii}) when there is no noise in the datasets. In Appendix~\\ref{sec:appb}, we noe that the null points in potential-field configurations are slightly affected by boundary conditions (periodic, closed, ...), spatial resolution or size of field-of-view: the null point with the strongest spectral radius remains stable.  From this study, we also provide a benchmark for analysing the topology of a magnetic configuration: it is possible to retrieve important information about the topology just by analysing the distribution of magnetic energy density and of the electric current density in the volume. Moreover, we emphasize the importance of checking the divergence-free property of the magnetic field in the vicinity of a null point.  \\begin{acks} SR thanks Eric Priest and Claire Parnell (University of St Andrews) for fruitful discussions on this topic. The computations of force-free field extrapolations have been performed using the {\\sf XTRAPOL} code developed by T. Amari (Ecole Polytechnique, France). \\end{acks}"
    },
    "1107/1107.1118_arXiv.txt": {
        "abstract": "{Studies of some bright, super-Eddington transient pulsars show a negative correlation between the energy of the cyclotron resonance scattering feature (CRSF) and the bolometric luminosity. For Her~X-1, using  repeated \\textsl{RXTE} observations during 1996--2005, the inverse dependence was found: the energy of the cyclotron line increases as the luminosity increases. The X-ray flux measured by the \\textsl{RXTE}/ASM (2--10\\,keV) has been assumed to represent the luminosity - more precisely: the maximum X-ray flux reached during the respective 35\\,d \\textsl{Main-On}. Here, we question whether the ASM flux is really an accurate measure of the bolometric luminosity of the source. We redetermined the energy of the cyclotron line and performed spectral fits using the combined data from the PCA (3.5--60\\,keV) and HEXTE (20--75\\,keV) instruments on \\textsl{RXTE} of the same 35\\,d cycles as used in the original work to determine the bolometric flux from those spectra. We confirm the result of the original analysis that the cyclotron line energy changes by $\\sim$7\\% for a change in flux by a factor of two. } ",
        "introduction": "Her X-1 is one of the most observed and well-studied accreting binary X-ray pulsars. Since the first observation made by  \\textsl{Uhuru}, the 35\\,d periodicity of the source is well known. The X-ray light curve shows two {\\em on}-states (high X-ray flux) and two {\\em off}-states (low X-ray flux), with  a \\textsl{Main-On} ($\\sim$ 7 orbital periods) and a \\textsl{Short-On} ($\\sim$ 5 orbital cycles), separated by two {\\em off}-states ($\\sim$ $4\\div5$ orbital cycles each). The maximum X-ray flux of the \\textsl{Main-on} is higher than the maximum X-ray flux of the \\textsl{Short-On} by a factor of three to four.  The modulation of the X-ray flux is believed to be caused by the periodic obscuration of the X-ray source by the precessing disk, which is believed to be inclined and warped \\citep{Gerend_Boynton_76}. The onset of the flux, the \\textsl{turn-on} (often identified with 35\\,d phase 0.0), is believed to occur when the outer rim of the disk opens up the view to the X-ray emitting regions near the polar caps on the surface of the neutron star, while the flux decrease towards the end of the {\\em on}-states is associated with the inner parts of the disk covering these regions from the observer.  Her X-1 is also the first X-ray pulsar  for which a cyclotron line was discovered \\citep{Truemper_etal_78}. The discovery of this line-like feature played a key-role in the measurement of the magnetic fields, providing the first direct measurement of the B-field of a neutron star. The energy of the cyclotron line is related to the magnetic field by the formula B$_{12}$ = (1+z) E$_{\\rm cyc}$/11.6\\,keV, where B$_{12}$ is the magnetic field strength in units of $10^{12}$\\,gauss, $z$ is the gravitational redshift and E$_{\\rm cyc}$ is the energy of the cyclotron line. This feature is now referred to as a cyclotron resonant scattering feature (CRSF) and seems to be quite common in accreting X-ray pulsars \\citep{Coburn_etal_02}. The cyclotron line is an absorption feature produced by the resonant scattering of photons on electrons. In the $\\sim10^{12}$\\,gauss magnetic field, the electrons are in quantized energy states (with respect to their movement perpendicular to the magnetic field), the so-called Landau levels. The energy levels are nearly equidistantly spaced and photons with energies equal to $n$ times the fundamental Landau energy could take part in this scattering.  In a few transient pulsars such as V0332+53 and 4U 0115+63, a negative correlation between the CRSF and the bolometric luminosity of the source has been observed:  the energy of the cyclotron line decreases as the X-ray luminosity increases \\citep{Mihara_etal_98,Mowlavi_etal_06,Nakajima_etal_06,Tsygankov_etal_06}. Her X-1, however, shows the opposite behavior. Repeated measurements of the cyclotron line energy with different instruments such as \\textsl{RXTE} and \\textsl{INTEGRAL} in the X-ray spectrum of Her X-1 have revealed  a \\textsl{positive} correlation between the (pulse phase-averaged) cyclotron line energy E$_{\\rm cyc}$ and the X-ray flux \\citep{Staubert_etal_07}, more precisely with the \\textsl{maximum X-ray flux} observed during the \\textsl{Main-On} of the respective 35\\,d cycle, as measured by \\textsl{RXTE}/ASM in the 2--10\\,keV range, which is assumed to be characteristic of the current accretion state and luminosity of the source. It has, however, been questioned whether the 2--10\\,keV flux can really be taken as representative of the bolometric luminosity.  In this paper, we attempt to answer this question by re-analysing the \\textsl{RXTE} observations of the same 35\\,d cycles as used in the original analysis of \\citet{Staubert_etal_07}. Using data from both instruments (PCA and HEXTE), we re-determine the cyclotron line energy and measure the 3.5--60\\,keV bolometric X-ray flux by performing spectral analyses of eleven \\textsl{Main-Ons} between 1996 and 2005. We show that the 2--10\\,keV flux is a good measure of the bolometric flux and that the positive correlation between the phase-averaged cyclotron line energy E$_{\\rm cyc}$ and the \\textsl{maximum 35\\,d flux}, or the luminosity, is confirmed.  \\begin{table} \\begin{center} \\vspace*{1mm} \\caption{Details of \\textsl{RXTE} observations of Her~X-1 used for the spectral analysis.} \\label{obs_table} \\resizebox{0.38\\textwidth}{!}{ \\begin{tabular}{l c c } \\hline\\hline Observation & 35~d Main-On              & Center  \\\\ month/year  & cycle number  $^{1}$    & MJD \\\\ \\hline July 96             &  257  & 50029.75 \\\\ September 97  &  269  & 50707.06 \\\\ December 00   &  303  & 51897.69 \\\\ January 01       &  304  & 51933.67 \\\\ May 01             &  307  & 52035.48 \\\\ June 01            &  308  & 52071.16 \\\\ August 02         &  320  & 52492.96 \\\\ November 02   &  323   & 52599.32 \\\\ December 02   &  324   & 52634.01\\\\ October 04       &  343  & 53300.95 \\\\ July 05              &  351  & 53577.35 \\\\ \\hline \\end{tabular} } \\vspace*{1mm} \\\\ $^{1}$    Cycle numbering according to \\citet{Staubert_etal_09} \\end{center} \\label{table1} \\end{table}  \\vspace*{-5mm} ",
        "conclusions": "In Fig. \\ref{compar} (left), we plot the newly determined cyclotron line energies against those from the original analysis. Overall, we find a very good agreement between the new values for E$_{\\rm cyc}$ and those from the original analysis: considering the respective uncertainties, the differences are smaller than 0.79 standard deviations for ten of the eleven values and  2.1 standard deviations for one value (cycle no. 323).   In Fig.~\\ref{compar} (right), the 3.5--60\\,keV maximum bolometric fluxes are plotted against the corresponding ASM fluxes of the original analysis. There is a good linear relationship between the two fluxes: the slope (taking the uncertainties in both variables into account) is 0.88$\\pm0.15$ ($\\rm 10^{-9}~erg~cm^{-2}~s^{-1}$)/(ASM cts s$^{-1}$), and the Pearson correlation coefficient is 0.82, corresponding to a chance probability of no correlation of P$<10^{-3}$. This demonstrates that the \\textsl{maximum ASM flux} can be taken as a good measure of the \\textsl{maximum bolometric X-ray flux} (and luminosity) of Her~X-1 during the respective 35\\,d cycle. The variation in flux from one 35\\,d cycle to the next is small, such that the maximum observed flux of a particular 35\\,d \\textsl{Main-On} can be considered a good measure of the luminosity of the source during this particular cycle. This is why the maximum ASM flux was used as a reference in the original analysis by \\citet{Staubert_etal_07}.  The final correlation between the cyclotron line energy and the X-ray flux is given in Fig.~\\ref{correl} in two ways: we correlate E$_{\\rm cyc}$ from the re-analysis with the \\textsl{scaled 3.5--60\\,keV maximum Main-On flux} (Fig.~\\ref{correl}, left), as well as the \\textsl{maximum ASM flux} (Fig.~\\ref{correl}, right). The corresponding slopes of the linear fits to these data (taking the uncertainties of both variables into account) and the corresponding Pearson correlation coefficients \\textsl{r} are: (i) for the Main-On flux presented in Fig.~\\ref{correl} (left) slope = 0.62$\\pm0.19$ (keV)/($\\rm 10^{-9}~erg~cm^{-2}~s^{-1}$) and \\textsl{r} = 0.62 (P=$2\\times10^{-2}$); (ii) for the maximum ASM flux presented in  Fig.~\\ref{correl} (right) slope = 0.67$\\pm0.14$ (keV)/(ASM cts s$^{-1}$) and \\textsl{r} = 0.90 (P=$10^{-4}$).  The correlation seen in Fig.~\\ref{correl} (left) is somewhat less convincing than that of Fig.~\\ref{correl} (right) (and that in the original analysis of \\citealt{Staubert_etal_07}). We attribute this to the unavoidably larger uncertainties associated mainly with the scaling of the bolometric flux measured for the individual spectrum to the flux that does describe the maximum flux of the particular 35\\,d cycle. The originally used \\textsl{maximum ASM fluxes} are, in contrast, simple direct measurements."
    },
    "1107/1107.6024_arXiv.txt": {
        "abstract": "Discerning the radiative dissipation mechanism for prompt emission in Gamma-Ray Bursts (GRBs) requires detailed spectroscopic modeling that straddles the $\\nu F_{\\nu}$ peak in the 100 keV - 1 MeV range.  Historically, empirical fits such as the popular Band function have been employed with considerable success in interpreting the observations. While extrapolations of the Band parameters can provide some physical insight into the emission mechanisms responsible for GRBs, these inferences do not provide a unique way of discerning between models. By fitting physical models directly this degeneracy can be broken, eliminating the need for empirical functions; our analysis here offers a first step in this direction. One of the oldest, and leading, theoretical ideas for the production of the prompt signal is the synchrotron shock model (SSM). Here we explore the applicability of this model to a bright {\\it Fermi} GBM burst with a simple temporal structure, GRB~{\\it 090820}A. Our investigation implements, for the first time, thermal and non-thermal synchrotron emissivities in the RMFIT forward-folding spectral analysis software often used in GBM burst studies. We find that these synchrotron emissivities, together with a blackbody shape, provide at least as good a match with the data as the Band GRB spectral fitting function.  This success is achieved in both time-integrated and time-resolved spectral fits. ",
        "introduction": "In the most popular paradigm for gamma-ray bursts of both long and short durations, it is typically assumed that prompt $\\gamma$-ray emission results from the dissipation of kinetic energy in a relativistically expanding fireball mediated by multiple internal shocks (e.g. see Piran 1999, or M\\'esz\\'aros 2001, for reviews). These shocks are presumed to diffusively accelerate a fraction of the electrons from thermal upstream distributions to higher energies. Usually only particles in the exponential tail of the Maxwellian are available for acceleration. Thus, for relativistic shocks, the expected outcome is that the particle distribution consists of a Maxwellian with a power-law tail at high energies. Based on this scenario, the radiative emission should consist of two components, quasi-thermal and non-thermal photons from electrons spiralling along magnetic field lines in optically thin regions of the jet. There could also be an additional photospheric contribution of Planckian form, originating in distinct, optically thick environs, perhaps interior to the regions spawning synchrotron emission.   To date, the characterization of GRB spectra has been dominated by the use of the empirical Band function \\citep{Band:1993p20784}, a parametrized, smoothly broken power law that was devised in the era of the BATSE experiment on the Compton Gamma-Ray Observatory (CGRO). Several authors have used measurements of the Band spectral shape parameters to infer properties of the physics involved in GRB emission. In particular, the fitted spectral indices defined by $N_{\\gamma}\\propto E^{-\\alpha}$ below the \\teq{\\nu F_{\\nu}} peak and $N_{\\gamma}\\propto E^{-\\beta}$ above it, may be compared with values predicted from synchrotron emission: the low-energy self-absorption index, $\\alpha$, of $+1$ (in photon flux units), the synchrotron `line of death'  index of $-2/3$, the `second line of death' at the fast cooling value of $-3/2$, the high energy index, $\\beta$, characterizing power-law particle acceleration, and the various spectral differences between these \\citep{Preece:1998p11011,Lloyd:2000p23470,LloydRonning:2002p17917,Preece:2002p11128}. However, it becomes difficult to discern between models through the Band function when the fitted low energy indices represent a power law only asymptotically, and when  many models predict similar Band indices. In fact, the Band function's inherent shape and curvature only loosely approximates the shape of the applicable physical models making it difficult to draw conclusions about emission mechanisms directly from Band function fits. A way to break this degeneracy is by fitting more realistic emission models to the data, which in addition provides deeper insights into the burst environment. In Section \\ref{sec:model} we detail the emission model that we use to fit GRB spectra. We present this model as a first step. In future work we will explore additional models in an attempt to discern between them. We describe our observational results with this model in Section \\ref{sec:observe}. ",
        "conclusions": "In this paper, we have shown that thermal and non-thermal synchrotron photon models, with an additional blackbody, are well consistent with the emission spectra of GRB {\\it 090820}A in various time intervals. These are physical models that afford the ability to constrain parameters that are physically meaningful, for example key descriptors of the electron distribution that is motivated by shock acceleration theory. By implementing these models into a forward-folding spectral analysis software we have been able to directly constrain many of the physical model parameters and their respective errors; a first in the field of GRB spectroscopy. This constitutes substantial progress over the use of the empirical Band function to fit prompt GRB spectra, which has been a nearly universal practice to date. The results presented here enable more rigorous statements about the validity of GRB emission models, moving the study of prompt burst emission into a new era.  Our modeling has focused on the standard synchrotron shock model with the addition of a blackbody component. The spectral fitting reveals a complex temporal evolution of the separate components. While spectral evolution is a well-known feature of GRBs, this type of fitting can enable \\textit{direct} physical interpretation of the evolution. These fits provide evidence that the line of death issue (Preece et al. 1998, 2002) can be overcome naturally with a combination of synchrotron and blackbody emission: the prominence of a blackbody component with its flat Rayleigh-Jeans portion would derive a comparably-fitted Band function with a flat low-energy index. This was also suggested by Guiriec et al. (2010) where the authors used simultaneous fits of the Band function and a blackbody. Note that it is possible that other physical models may, in fact, produce superior fits to the data for GRB~{\\it 090820}A and other bursts. Strongly-cooled synchrotron emission, inverse Compton and jitter radiation are popular candidates, and our work here motivates the future development of RMFIT software modules for these processes.      A principal finding of the analysis in this paper is that the power-law synchrotron component is orders of magnitude more intense than the thermal synchrotron component during the peak of the burst, the latter contributing at most a few percent of the flux. This confirms the finding of BB04 for BATSE/EGRET bursts GRB {\\it 910503}, GRB {\\it 910601} and GRB {\\it 910814}, which was a theoretically-based perspective that did not fold models through the detector response matrices. They had noted that full plasma and Monte Carlo diffusion simulations of shock acceleration clearly predict a power-law tail in the particle distribution that smoothly extends from the dominant thermal population (e.g. see also Baring {2011}, and references therein). This tail is several orders of magnitude smaller than what is found when fitting synchrotron emission to burst spectra. It is not clear how such non-thermally-dominated distributions can arise near shocks, providing a conundrum for the standard synchrotron shock model. Limited smoothing of the sharp peak of the non-thermal electron component will not alter this conclusion.   This result is also in accord with Guiriec et al. (2010), in their analysis of GRB~{\\it 100724}B, who fitted its GBM spectra with a combination of the Band model and a blackbody. They too found that an unrealistically high efficiency for the acceleration mechanism or a source size smaller than the innermost stable orbit of a black hole was required to invoke the standard fireball model for explaining the origin of the $\\gamma$-ray emission. Therefore, it was surmised therein that the outflow from the jet was at least partially magnetized.   To conclude, the success of this analysis in isolating the relative contributions of a handful of distinct spectral components indicates that it is imperative for the field of GRB spectroscopy to move away from the use of the empirical fitting functions: many physical models can asymptotically approximate the Band spectral indices, rendering it difficult to discern between them particularly near the $\\nu F_{\\nu}$ peak. Instead, direct comparisons of the fitted physical models are possible, and are required to truly discriminate between the various emission processes. The fitting of physical SSM/blackbody spectra here offers a clear advance beyond empirical fits, and provides the impetus for further development and deployment of physical modeling of prompt burst emission spectra."
    },
    "1107/1107.0508_arXiv.txt": {
        "abstract": "RX J1713.7$-$3946 is the most remarkable TeV $\\gamma$-ray SNR which emits $\\gamma$-rays in the highest energy range. We made a new combined analysis of CO and \\ion{H}{1} in the SNR and derived the total protons in the interstellar medium (ISM). We have found that the inclusion of the \\ion{H}{1} gas provides a significantly better spatial match between the TeV $\\gamma$-rays and ISM protons than the H$_2$ gas alone. In particular, the southeastern rim of the $\\gamma$-ray shell has a counterpart only in the \\ion{H}{1}. The finding shows that the ISM proton distribution is consistent with the hadronic scenario that comic ray (CR) protons react with ISM protons to produce the $\\gamma$-rays. This provides another step forward for the hadronic origin of the $\\gamma$-rays by offering one of the necessary conditions missing in the previous hadronic interpretations. We argue that the highly inhomogeneous distribution of the ISM protons is crucial in the origin of the $\\gamma$-rays. Most of the neutral gas was likely swept up by the stellar wind of an OB star prior to the SNe to form a low-density cavity and a swept-up dense wall. The cavity explains the low-density site where the diffusive shock acceleration of charged particles takes place with suppressed thermal X-rays, whereas the CR protons can reach the target protons in the wall to produce the $\\gamma$-rays. The present finding allows us to estimate the total CR proton energy to be $\\sim $10$^{48}$ ergs, 0.1 \\% of the total energy of a SNe. ",
        "introduction": "It is a long-standing question how the cosmic ray (CR) protons, the major constituent of cosmic rays, are accelerated in the interstellar space. Supernova remnants (SNRs) are the most likely candidate for the acceleration because the high-speed shock waves offer an ideal site for diffusive shock acceleration (DSA) \\citep[e.g.,][]{bell1978, blandford1978}. The principal site of CR proton acceleration is, however, not yet identified observationally in spite of a number of efforts to address this issue.  RX J1713.7$-$3946 is the brightest and most energetic TeV $\\gamma$-ray SNR detected in the Galactic plane survey with H.E.S.S. \\citep{aharonian2006a}  and is a primary candidate where the origin of the $\\gamma$-rays may be established. Discovery of the SNR was made in X-rays with {\\it ROSAT} \\citep{pfeffermann1996} and {\\it ASCA} showed that the X-rays are non-thermal synchrotron emission with no thermal features \\citep{koyama1997}. TeV $\\gamma$-rays were first detected by CANGAROO \\citep{enomoto2002} and, subsequently, H.E.S.S. resolved the shell-like TeV $\\gamma$-ray distribution with a $\\sim 0\\fdg 1$ degree point spread function \\citep{aharonian2004, aharonian2006b, aharonian2007}. The $\\gamma$-rays are emitted via two mechanisms, either leptonic or hadronic, closely connected to the CR particles and it is important to understand which mechanism is working in the SNR. The leptonic process includes the inverse Compton effect of CR electrons which energize the low energy photons of the cosmic microwave background and additional soft photon fields (e.g., infrared). The hadronic process includes the neutral pion decay into $\\gamma$-rays following the reaction of CR protons with the low energy target protons in the interstellar medium (ISM). Considerable work has been devoted to explaining the $\\gamma$-ray emission in the framework of leptonic and hadronic scenarios \\citep{aharonian2006b, porter2006, katz2008, berezhko2008, ellison2008, tanaka2008, morlino2009, acero2009, ellison2010, patnaude2010, zirakashvili2010, abdo2011, fang2011}.  The molecular gas interacting with the SNR was discovered in $V_{\\mathrm{LSR}}$, the velocity with respect to the local standard of rest, around $-7$ km s$^{-1}$ at 2.6 arcmin resolution based on NANTEN Galactic plane CO survey and the distance of the SNR was determined to be 1 kpc by using the flat rotation curve of the Galaxy \\citep{fukui2003}. This determination offered a robust verification for a small distance of 1 kpc and revised an old value 6 kpc derived from lower resolution CO observations \\citep{slane1999}. Studies of X-ray absorption suggested a similar distance 1 kpc under an assumption of uniform foreground gas distribution \\citep{koyama1997}, but the local bubble of \\ion{H}{1} located by chance toward the SNR makes the X-ray absorption uncertain in estimating the distance \\citep{slane1999, matsunaga2001}. A subsequent careful analysis of the X-ray absorption favors the smaller distance \\citep{cassamchenai2004}. At 1 kpc the SNR has a radius of 9 pc and an age of 1600 yrs \\citep{fukui2003, wang1997} and the expanding shock front has a speed of 3000 km s$^{-1}$ \\citep{zirakashvili2007, uchiyama2003, uchiyama2007}. \\citet{moriguchi2005}  showed further details of the NANTEN CO distribution and confirmed the identification of the interacting molecular gas at $-20$--0 km s$^{-1}$ by \\citet{fukui2003}. Most recently, \\citet{sano2010}  showed that the SNR harbors the star forming dense cloud core named peak C at similar $V_{\\mathrm{LSR}}$ and argued that X-rays are bright around the core, reinforcing the association of the molecular gas.  Importantly, the associated molecular gas with the SNR opened a possibility to identify target protons where the hadronic process is working. If the cosmic ray density is nearly uniform, we expect the $\\gamma$-ray distribution mimics that of the interstellar target protons. Some nearby molecular clouds show good spatial correlation with relatively high resolution $\\gamma$-ray images, clearly verifying that the hadronic process is working to produce $\\gamma$-rays in the cosmic ray sea (e.g., \\citeauthor{bertsch1993} \\citeyear{bertsch1993} and the references therein.). A detailed comparison between the ISM protons and the recent high-resolution $\\gamma$-ray images of H.E.S.S. sources is useful to test the correlation, although such a test was not possible until recently in the preceding low resolution $\\gamma$-ray observations at degree-scale resolution. \\citet{aharonian2006b} compared the NANTEN CO distribution with the H.E.S.S. TeV $\\gamma$-ray image and examined both leptonic and hadronic scenarios as the origin of the $\\gamma$-rays. By adopting annular averaging of TeV $\\gamma$-rays and CO in the shell \\citep[see Figure 17 of][]{aharonian2006b}, these authors found that TeV $\\gamma$-rays are fairly well correlated with CO, whereas the correlation is not complete in the sense that the southeastern rim of the TeV $\\gamma$-ray shell has no counterpart in CO. The complete identification of target ISM protons thus remained unsettled in the hadronic scenario.  The $\\gamma$-rays and X-rays are significantly enhanced toward the clumpy molecular gas at a pc scale as first shown by \\citet{fukui2003}. A strong connection among the molecular gas, the $\\gamma$-rays, the X-rays, and perhaps cosmic rays is therefore suggested \\citep{fukui2003, fukui2008, sano2010, zirakashvili2010}. Most of the previous models of the $\\gamma$- and X-rays cited earlier assume more or less uniform density distribution of the ISM in the SNR, whereas the observations of the ISM indicate that the actual distribution is highly inhomogeneous, with density varying by a factor 100 or more around the SNR. In addition, Galactic-scale studies of $\\gamma$-rays suggest that there is \"dark gas\" which is not detectable in CO or in \\ion{H}{1} but still contributes to the $\\gamma$-rays and visual extinction \\citep{grenier2005, ade2011}. Such gas may be either cold \\ion{H}{1} or H$_2$ with no detectable CO. So, it is important to consider \\ion{H}{1} carefully in order to have a comprehensive understanding of the ISM protons. In the present paper, we shall use the term \"dark \\ion{H}{1}\" for observed \\ion{H}{1} with significantly lower brightness than the surroundings. So, \"dark \\ion{H}{1}\" does not necessarily mean \"dark gas\" above, while they might be linked.  We here present a combined analysis of both the $^{12}$CO($J$=1--0) and \\ion{H}{1} datasets in order to clarify the distribution of the ISM protons in RX J1713.7$-$3946. The quantitative analysis is made mainly by using the $^{12}$CO($J$=1--0) and \\ion{H}{1} data here and a comparative study of the $^{12}$CO($J$=1--0, 2--1) transitions is a subject of a forthcoming paper. A theoretical work of magneto-hydrodynamics which incorporates fully the inhomogeneous ISM in RX J1713.7$-$3946 is complimentary to the present observational paper and will be published separately (Inoue, Yamazaki, Inutsuka $\\&$ Fukui 2011; hereafter IYIF2011). The present paper is organized as follows. Section \\ref{section:datasets} gives the description of the CO and \\ion{H}{1} datasets. Section \\ref{section:analysis} consists of five Sub-sections; Sub-section 3.1 gives a combined data set of CO and \\ion{H}{1}, Sub-sections 3.2--3.4 the distribution of total ISM protons in the SNR with an emphasis on dark \\ion{H}{1}, the cold and dense atomic phase of ISM protons, and Sub-section 3.5 comparison with the $\\gamma$-rays. Section \\ref{section:discussion} consists of two Sub-sections; Sub-section 4.1 describes the initial distribution of the ISM before the stellar explosion, and Sub-section 4.2 the scheme of particle acceleration and its relationship with the $\\gamma$-rays. The conclusions are given in Section \\ref{section:conclusions}.   \\begin{figure*} \\begin{center} \\includegraphics[width=179mm,clip]{f1.eps} \\caption{(a) The H.E.S.S. TeV $\\gamma$-ray distribution of RX J1713.7$-$3946 in smoothed excess counts above the cosmic-ray background (see Figure 2 of \\citeauthor{aharonian2007} \\citeyear{aharonian2007}). Contours are plotted every 10 smoothed counts from 20 smoothed counts. (b) Averaged brightness temperature distribution of $^{12}$CO($J$=1--0) emission in a velocity range of $V_{\\mathrm{LSR}}$ = $-20$ km s$^{-1}$ to $0$ km s$^{-1}$ is shown in color \\citep{fukui2003, moriguchi2005}. White contours show the H.E.S.S. TeV $\\gamma$-ray distribution and are plotted every 20 smoothed counts from 20 smoothed counts. (c) Averaged brightness temperature distribution of \\ion{H}{1} emission obtained by ATCA and Parkes in a velocity range from $V_{\\mathrm{LSR}}$ = $-8$ km s$^{-1}$ to $-6$ km s$^{-1}$ \\citep{mccluregriffiths2005} is shown in color. White contours show the $^{12}$CO($J$=1--0) brightness temperature integrated in the same velocity range every 1.0 K km s$^{-1}$ ($\\sim$3$\\sigma$).} \\end{center} \\end{figure*}% ",
        "conclusions": "\\label{section:conclusions} We summarize the main conclusions as follows;  \\begin{enumerate} \\item A new analysis of CO and \\ion{H}{1} has revealed that the TeV $\\gamma$-ray SNR RX J1713.7$-$3946 is associated with a significant amount of \\ion{H}{1} gas without H$_2$ derived from CO. This \\ion{H}{1} gas is relatively dense and cold and detectable mainly as \\ion{H}{1} emission. We have also identified regions where \\ion{H}{1} is observed as dark \\ion{H}{1} in self-absorption dips and derived the total ISM proton column density over the SNR. The \\ion{H}{1} plus H$_2$, the total ISM protons, provides one of the necessary conditions, target protons,  in the hadronic origin of the $\\gamma$-rays. Such target ISM protons have not been identified in the previous study that took into account only H$_2$, although the present finding alone does not exclude the leptonic origin.  \\item For an annular pattern around the TeV $\\gamma$-ray shell, we compared the total ISM proton distribution with the TeV $\\gamma$-ray distribution and found that they show reasonably good correspondence, varying by similar factors. The inclusion of the atomic protons observed as the \\ion{H}{1} self-absorption dips is essential particularly in the southeast of the $\\gamma$-ray shell. The interpretation of \\ion{H}{1} self-absorption dips is also supported by the enhanced optical extinction toward the southeast rim.  \\item The cavity surrounding the SNR was created by the stellar wind of the SN projenitor. The inside of the cavity is of low density with $<$1 cm$^{-3}$ while the cavity wall consists of the dense and clumpy atomic or molecular target protons of $\\geq $100--1000 cm$^{-3}$. The diffusive shock acceleration in the highly inhomogeneous ISM offers a reasonable mechanism of particle acceleration in the low-density cavity and the dense wall acts as the target for $\\gamma$-ray production by the CR protons. Hydro-dynamical numerical simulations of the interaction have shown detailed physical processes involved (IYIF2011).  \\item By considering the other pieces of the observational and theoretical works accumulated thus far, the present results make the hadronic interpretation much more comfortable in RX J1713.7$-$3946. The current energy of the total CR protons is estimated to be $\\sim $10$^{48}$ ergs, 0.1 $\\%$ of the total energy of SNe, if we assume the $\\gamma$-rays are all produced by the hadronic process. \\end{enumerate}"
    },
    "1107/1107.0813_arXiv.txt": {
        "abstract": "We study a multi-field inflationary theory with separable Lagrangian, which has different speed of sound for each field. We find that the fields always coupled at perturbative level through gravitational interaction. We show that if the coupling terms among the perturbation fields are weak enough, these fields can be treated as a combination of decoupled fields, which are similar to normal modes in coupled oscillation. By virtue of such fields, the curvature perturbation at the horizon-crossing can be calculated up to the leading order of slow variation parameters via $\\delta \\mathcal{N}$ formalism. Explicitly, we consider a model of multi-speed DBI inflation, and calculate the power spectrum in detail. The result depends on the ratio of different speeds of sound, and shows an apparent amplification when the ratio deviates from unity. ",
        "introduction": "Inflationary cosmology has become the prevalent paradigm to understand the early stage of our universe, with its advantages of resolving the flatness, homogeneity and monopole problems \\cite{inflation_bible, Starob}, and predicting a approximately scale-invariant primordial power spectrum consistent with current cosmological observations \\cite{Komatsu:2010fb} very well. However, a single field inflation model often suffers from fine tuning problems on the parameters of its potential, such as the mass and the coupling constant.  In recent years, people has noticed that, when a number of scalar fields are involved in the inflationary stage, they can relax many limits on the single scalar inflation model \\cite{Liddle:1998jc}. Usually, these fields are able to work cooperatively to give an inflationary stage long enough, even none of them can sustain inflation separately. Models of this type have been considered later in Refs. \\cite{Malik:1998gy, Kanti:1999vt, Copeland:1999cs, Green:1999vv}. The main results showed that both the e-folding number $\\mathcal{N}_e$ and the curvature perturbation $\\mathcal{R}$ are approximately proportional to the number of the scalars $N$. Later, the model of N-flation was proposed by Dimopoulos {\\it et al.} \\cite{Dimopoulos:2005ac}, which showed that a number of axions predicted by string theory can give rise to a radiatively stable inflation. This model has explored the possibility for an attractive embedding of multi-field inflation in string theory.  Over the past several years, based on the recent developments in string theory, there have been many studies on its applications to the early universe in inflationary cosmology. An interesting inflation model, which has a non-canonical kinetic term inspired by string theory, was studied intensively in the literature. Due to a non-canonical kinetic term, the propagation of field fluctuations in this model is characterized by a sound speed parameter and the perturbations get freezed not on Hubble radius, but on the sound horizon instead. One specific realization of this type of models can be described by a Dirac-Born-Infeld-like (DBI) action \\cite{Aharony:1999ti, Myers:1999ps}. Based on brane inflation \\cite{Dvali:1998pa}, the model with a single DBI field was investigated in detail\\cite{Silverstein:2003hf, Chen:2004gc, Chen:2005ad}, which has explored a window of inflation models without flat potentials. In this model, a warping factor was applied to provide a speed limit which keeps the inflaton near the top of a potential even if the potential is steep.  Motivated by the effective description of multiple D-brane dynamics in string theory \\cite{Taylor:1999pr}, an interesting inflation model involving multiple sound speeds with each sound speed characterizing one field fluctuation was proposed \\cite{Cai:2009hw, Cai:2008if,Christopherson:2008ry}. The authors of Refs.\\cite{Cai:2009hw, Cai:2008if} suggest this scenario can be realized by a number of general scalar fields with arbitrary kinetic forms, and these scalars have their own sound speeds respectively. Therefore, this model is dubbed as ``{\\it Multi-Speed Inflation}\" (MSI). In this model, the propagations of field fluctuations are individual, and the usual conceptions in multi-field inflation models might be not suitable in this scenario. For example, in a usual generalized N-flation model, the length scale for perturbations being freezed takes the unique sound horizon; however, in our model it corresponds to the maximum sound horizon. It is worth emphasizing that the model of MSI is different from the usual DBI N-flation in which only multiple moduli fields are involved in one DBI action \\cite{Ward:2007sj, Easson:2007fz, Huang:2007hh, Ward:2007gs, Langlois:2008mn, Contaldi:2008hr, Langlois:2009ej}, but MSI is constructed by multiple Kessence type actions. For example, an explicit model of MSI is made of two DBI fields in Ref \\cite{Cai:2009hw}, and its primordial perturbations including curvature and entropy modes and their non-Gaussianities were considered.  In the current paper we extend the model into a much more generic case by relaxing the form of the Lagrangian. We find that even there are no coupling terms among the inflaton fields, their perturbation modes are coupled due to the nonlinear gravitational interaction. We also find that if these couplings are weak, the independent components of the fields can be treated as a combination of normal modes, by which the curvature perturbation at the horizon-crossing can be calculated in detail up to leading order of slow variation parameters via $\\delta \\mathcal{N}$ formalism. As a specific example, we consider a model of multi-speed inflation involving two DBI fields. We show that in the relativistic limit, the coupling between two fields is mainly reflected by the damping terms in the perturbation equations. The difference between the sound speeds of these two fields is able to generate a considerable amplification on the amplitude of the primordial curvature perturbation. This is greatly different from the usual analysis on the inflationary model constructed by two canonical fields\\cite{Wands:2002bn, Gordon:2000hv}.   The paper is organized as follows. In Section II, we study the cosmological perturbation theory of the MSI model in a generic case, then show that the inflaton fluctuations could be decoupled through a redefinition of fields at leading order under the assumption of weak couplings. In Section III, we focus on a model of MSI involving two fields, and study the detailed field transformation matrix to illustrate that the decoupling process of the field fluctuations is reliable under the weak coupling approximation. Specifically, we analyze a specific model constructed by two DBI fields and give its curvature perturbation. Section IV presents a summary and discussion.  In the paper we take the normalization $M_p^2=1/8\\pi G=1$ and the sign of metric is adopted as $(-,+,+,+)$ in the following. ",
        "conclusions": "In the frame of inflationary cosmology involving multiple inflaton fields, it is possible to allow these fields possess their own sound speed parameters which characterize the propagation of field fluctuations correspondingly. It is important to study how these modes are related to the cosmological perturbation we may observe in CMB experiments. In the present paper we analyzed the dynamics of field fluctuations in a general scenario of MSI. Assuming no direct coupling terms among the inflaton fields, we noticed that these perturbation modes are generically coupled both through their effective mass terms and cosmic damping terms. Fortunately, we found that if these couplings are weak, the perturbation theory can be treated as a combination of normal modes, by which the curvature perturbation at the horizon-crossing can be calculated in detail up to leading order of slow-variation parameters via $\\delta \\mathcal{N}$ formalism for a specific model.  Specifically we studied a model of multi-speed inflation involving two DBI fields with their sound speeds being small and slow-varying. Our analysis showed that in the relativistic limit, the coupling between two field fluctuations is mainly contributed by the damping terms in their perturbation equations. We further studied the curvature perturbation in this model, and verified their primordial power spectra are nearly scale-invariant. This has interesting implications to the curvaton scenario\\cite{Lyth:2001nq}, which suggests a light field in inflationary phase could seed entropy perturbation and convert it into curvature perturbation at the end of inflation\\cite{Mollerach:1989hu, Linde:1996gt, Enqvist:2001zp, Lyth:2001nq, Moroi:2001ct}\\footnote{Also see Ref. \\cite{Cai:2011zx} for the extension of the curvaton scenario to non-inflationary cosmology.}. Namely, as shown in \\cite{Li:2008fma, Zhang:2009gw, Cai:2010rt}, a light DBI field can realize the curvaton scenario and generate sizable non-Gaussianities of local and equilateral types after suitable fine-tuning on the model parameters. Our study also showed that, when the sound speed parameters for the inflaton fields are not equal, the perturbation modes would get frozen at different sound horizons. An important phenomenon is that the curvature perturbation of MSI could obtain an enhancement which depends on the ratio of the sound speeds.   The scenario of MSI is a very important branch of various models of multiple field inflation. In recent years, the inflationary models involving multiple field components have been studied extensively in the literature, namely, the analysis on the dynamics of N-flation \\cite{Kim:2006ys}; multiple field inflation with particle decays \\cite{Choi:2007fya}; the N-flation model in the frame of braneworld \\cite{Panotopoulos:2007pg}; the model of staggered Nflation in the stringy landscape \\cite{Battefeld:2008py}; dynamics of the model of coupled N-flation with canonical fields \\cite{Ashoorioon:2008qr}; the inflation model constructed by a DBI field coupled to radiation \\cite{Cai:2010wt}; and see \\cite{Wands:2007bd, Langlois:2008sg} for comprehensive reviews on this field. It has been well realized that, for the model of multi-field inflation, when the sound speeds for the fields are the same, the field fluctuations are able to be rotated orthogonally to a basis in which the curvature and entropy fluctuations decouple explicitly, at least up to leading order. However, when we relax the scenario by allowing the sound speed parameters are different, our analysis implied that the rotation of the field fluctuations are no longer orthogonal but we still are able to find a transformation to decouple the fields under some conditions and obtain the curvature perturbation via $\\delta\\mathcal{N}$ formalism."
    },
    "1107/1107.0678_arXiv.txt": {
        "abstract": "According to recent models, the accretion disk and black hole in active galactic nuclei are surrounded by a clumpy torus. We investigate the NIR flux variation of the torus in response to a UV flash for various geometries. Anisotropic illumination by the disk and the torus self-occultation contrast our study with earlier works. Both the waning effect of each clump and the torus self-occultation selectively reduce the emission from  the region with a short delay. Therefore, the NIR delay depends on the viewing angle (where a more inclined angle leads to a longer delay) and the time response shows an asymmetric profile with a negative skewness, opposing to the results for optically thin tori. The range of the computed delay coincides with the observed one, suggesting that the viewing angle is primarily responsible for the scatter of the observed delay. We also propose that the red NIR-to-optical color of type-1.8/1.9 objects is caused by not only the dust extinction but also the intrinsically red color. Compared with the modest torus thickness, both a thick and a thin tori display the weaker NIR emission. A selection bias is thus expected such that NIR-selected AGNs tend to possess moderately thick tori. A thicker torus shows a narrower and more heavily skewed time profile, while a thin torus produces a rapid response. A super-Eddington accretion rate leads to a much weaker NIR emission due to the disk self-occultation and the disk truncation by the self-gravity. A long delay is expected from an optically thin and/or a largely misaligned torus. A very weak NIR emission, such as in hot-dust-poor active nuclei, can arise from a geometrically thin torus, a super-Eddington accretion rate or a slightly misaligned torus. ",
        "introduction": "Active galactic nuclei (AGNs) are powered by gas accretion onto supermassive black holes (BHs) at the center of each galaxy. A variety of observations suggest that the accretion disk and the BH are surrounded by an optically and geometrically thick torus (Telesco et al. 1984; Antonucci \\& Miller 1985; Miller \\& Goodrich 1990; Radovich et al. 1999). Since the torus potentially plays a role of a gas reservoir for the accretion disk, its nature, such as the structure, the size and the mass, has long been investigated (Pier \\& Krolik 1992, 1993; Fukue \\& Sanbuichi 1993; Granato \\& Danese 1994; Efstathiou \\& Rowan-Robinson 1995 Beckert \\& Duschl 2004; Mor et al. 2009).  A large geometrical thickness of the torus revealed by various observations (Antonucci 1993; Pogge 1989; Wilson \\& Tsvetanov 1994; Schmitt \\& Kinney 1996) indicates that numerous dusty clumps, rather than a smooth mixture of gas and dust, constitute the torus with a large clump-to-clump velocity dispersion $\\sim \\! 100\\,\\mathrm{km\\, s}^{-1}$ (Krolik \\& Begelman 1988; Wada \\& Norman 2002; H\\\"{o}nig \\& Beckert 2007). Temperature of clumps is less than a critical temperature $T_{\\rm sub} \\sim 1500$\\,K above which dust grains are sublimated (Barvainis 1987). Infrared (IR) emission and absorption features provide unique opportunities to probe the clumpy torus (Nenkova et al. 2002, 2008; Dullemond \\& van Bemmel 2005; H\\\"{o}nig et al. 2006; Geballe et al. 2006; Shirahata et al. 2007; Ibar \\& Lira 2007; Schartmann et al. 2008; Deo et al. 2011).   Clumps are heated by illumination from the central accretion disk, and closer clumps have higher temperature. Inner edge of the dusty torus is determined by the sublimation process so that clumps' temperature equals $T_{\\rm sub}$ there, and radiates at Near-IR (NIR) % as \"3$\\mu$m bump\" (Rees et al. 1969; Neugebauer et al. 1979; Edelson \\& Malkan 1986; Kobayashi et al. 1993). Based on the energy balance of the clump closest to the BH, Barvainis (1987) derived the innermost radius of the torus (dust sublimation radius, denoted as $R_{\\rm sub,0}$ in this study): \\begin{equation} R_{\\rm sub,0} = 0.13 \\brfrac{L_{\\rm UV}}{10^{44} \\, \\mathrm{erg}\\,\\mathrm{s}^{-1}}^{0.5} \\brfrac{T_{\\rm sub}}{1500 \\,\\mathrm{K}}^{-2.8} \\brfrac{a}{0.05 \\,\\mu \\mathrm{m}}^{-0.5} {\\rm pc}, \\label{eq:bar87} \\end{equation} where $L_{\\rm UV}$ and $a$ are UV luminosity and the size of dust grains, respectively.  Indeed, NIR emission from type-1 AGNs lags behind optical variation by an order of a month (Clavel et al. 1989; Glass 1992, 2004; Nelson 1996; Oknyanskij et al. 1999; Minezaki et al. 2004; Suganuma et al. 2004). Moreover, the luminosity dependency of the time lag also coincides with the theoretical prediction as $\\propto L_{\\rm UV}^{0.5}$ (Suganuma et al. 2006; Gaskell et al. 2007). However, the NIR-to-optical time lag is systematically smaller than the lag predicted from Equation (\\ref{eq:bar87}) by a factor of $\\sim \\! 1/3$ (Oknyanskij \\& Horne 2001; Kishimoto et al. 2007; Nenkova et al. 2008). To tackle with this conflict, Kawaguchi \\& Mori (2010, hereafter Paper I) pointed out that the illumination by an optically thick disk is inevitably anisotropic, which is a fact missing in deriving Equation (\\ref{eq:bar87}). There is a systematic difference between the inclination angle at which we observe the disk in type-1 AGNs and the angle at which an aligned torus observes the disk. The effects of the anisotropic % illumination naturally resolve the puzzle of the systematic deviation of a factor of $\\sim \\! 1/3$ (Paper I).  In Paper I, we assumed the configuration appropriate for a typical type-1 AGN. In this study, we investigate the expected characteristics of the NIR emission for various geometries of the disk, the torus and the observer. Anisotropic illumination by the disk and the effect of the torus self-occultation contrast our study with earlier works. The next section describes the calculational methods of our model. Then, properties of NIR emission from an aligned (Section \\ref{sec:aligned}) and a misaligned (Section \\ref{sec:misaligned}) tori are presented. Finally, we make a summary % of this study in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  According to recent models, the accretion disk and the BH in AGNs are surrounded by a clumpy torus, with its inner radius governed by the dust sublimation process. Regarding the inner radius of the torus, there was a systematic deviation between the observational results and the theory. In Paper I, we showed that the anisotropy of the disk emission resolves this conflict for % a typical type-1 AGN. We found that the anisotropy makes the torus inner region closer to the central BH and concave. Furthermore, the innermost edge of the torus may connect with the outermost edge of the accretion disk continuously.  In this study, we have calculated the NIR flux variation of the torus in response to a UV flash for various geometries of the disk, the torus and the observer. Anisotropic illumination by the disk and the effect of the torus self-occultation contrast our study with earlier works. We have found that both the waning effect of each clump and the torus self-occultation selectively reduce the emission from  the region with a short delay. Thus, the resultant NIR time response shows a $\\theta_{\\rm obs}$-dependent delay and an asymmetric profile with a negative skewness, opposing to the results for optically thin tori (Barvainis 1992).  By contrast with the fiducial viewing angle % of 25\\degr, a small viewing angle results in a short time delay with a narrow and peaky response. On the other hand, a more inclined viewing angle leads to a longer delay with a broader profile and to % a redder NIR-to-optical color. We propose that the red NIR-optical color of type-1.8/1.9 objects is caused by not only the dust extinction but also intrinsically red color. The computed range of $t_{\\rm delay}$ coincides with the observed one.  Compared with the modest torus thickness of $45\\degr$, both a thick and a thin tori display the weaker NIR emission, consistent with the work by Nenkova et al. (2008). In other words, a modest thickness of the torus leads to the strongest NIR emission. % A selection bias is thus expected such that NIR-selected AGNs tend to possess moderately thick tori. For thin tori, % we have found that the NIR fluence is in proportion to $\\Omega_{\\rm torus}^{1.9}$. As the torus becomes thicker, the NIR response shows a slightly longer delay with a narrower and more heavily skewed profile due to the torus self-occultation. % This trend opposes to the % viewing angle dependency, where the delay and the width are positively correlated each other. On the contrary, as the torus gets thinner, the NIR response becomes more rapid, narrower, closer to time-symmetric and low. %  In contrast to % a thin disk with a sub-Eddington accretion rate, % a super-Eddington accretion rate leads to a much weaker ($< \\! 1/5$) NIR emission due to the disk self-occultation and the disk truncation by the self-gravity, and to a low % time response.  Among the three dependencies examined for aligned tori, only the variation of the viewing angles reproduces the observed range of the delay. Therefore, we propose that the viewing angle is the key parameter responsible for the observed scatter about the regression line in the $t_{\\rm delay}$ v.s. $L_{\\rm UV}$ diagram. Conversely, the measurements of $t_{\\rm delay}$ are potentially useful to estimate the inclination angles.  We have also investigated the consequences of the misalignment between the torus and the disk axes. A variety of $t_{\\rm delay}$ is achieved, which is wide enough to cover the observed range. A short delay % is obtained for a small misalignment and is associated with a small NIR fluence, while a % a long delay is obtained for a largely misaligned torus with an usual fluence. % This trend contrasts with the % viewing angle dependency for aligned tori, where a short delay is associated % with a normal NIR fluence while a long delay % means a large NIR fluence. For highly misaligned cases, contributions from $\\theta > \\frac{\\pi}{2}$ increase the time delay (up to $\\sim R_{\\rm sub,0}/c$), the NIR fluence and the skewness ($s \\ga 0$).  In case the torus is optically thin (with an inefficient self-occultation), the time delay of the NIR emission from an aligned torus becomes longer. Moreover, misaligned optically thin tori % show a yet longer delay ($> R_{\\rm sub,0}/c$) with a huge NIR flux.  From an observational point of view, these numerical results are summarised as follows. If the observed time delay of the NIR emission is short, it will mean either a small viewing angle, a geometrically thin torus, or a slightly misaligned torus. If the delay is long, on the other hand, it indicates either an inclined viewing angle, an aligned optically thin torus, or a largely misaligned torus. An extremely long delay ($> R_{\\rm sub,0}/c$) would mean a misaligned optically thin torus. As to the NIR flux, a blue NIR-to-optical color (i.e., a weak NIR emission such as in hot-dust-poor AGNs in the weakest cases) indicates either a geometrically thin torus, a geometrically thick torus, a super-Eddington accretion, or a slight misalignment between the torus and the disk. On the other hand, a red NIR-optical color (a large NIR emission) means a large viewing angle with a modest geometrical thickness of the torus or a largely misaligned optically thin torus."
    },
    "1107/1107.0921.txt": {
        "abstract": "The kinematical evolution of four EUV waves, well observed by the Extreme UltraViolet Imager (EUVI) onboard the Solar-Terrestrial Relations Observatory (STEREO), is studied by visually tracking the wave fronts as well as by a semi-automatized perturbation profile method leading to results matching each other within the error limits. The derived mean velocities of the events under study lie in the range of 220--350~km\\,s$^{-1}$. The fastest of the events (May 19, 2007) reveals a significant deceleration of $\\approx -190$~m~s$^{-2}$ while the others are consistent with a constant velocity during the wave propagation. The evolution of the maximum intensity values reveals initial intensification by 20 up to 70\\%, and decays to original levels within 40--60~min, while the width at half maximum and full maximum of the perturbation profiles are broadening by a factor of 2\\,--\\,4. The integral below the perturbation profile remains basically constant in two cases, while it shows a decrease by a factor of 3\\,--\\,4 in the other two cases. From the peak perturbation amplitudes we estimate the corresponding magneto-sonic Mach numbers M$_{ms}$ which are in the range of 1.08--1.21. The perturbation profiles reveal three distinct features behind the propagating wave fronts: coronal dimmings, stationary brightenings and rarefaction regions. All of them appear after the wave passage and are only slowly fading away. Our findings indicate that the events under study are weak shock fast-mode MHD waves initiated by the CME lateral expansion. ",
        "introduction": "Large-scale, large-amplitude disturbances propagating through the solar corona have been first observed by \\citet{moses97} and \\citet{thompson98} in images recorded by the Extreme Ultraviolet Imaging Telescope \\citep[EIT;][]{delaboudiniere95} onboard the Solar and Heliospheric Observatory \\citep[SoHO;][]{domingo95}, thereafter called `EIT waves'. They were originally interpreted as the coronal counterparts of the chromospheric Moreton waves \\citep{moretonramsey60} as suggested by \\citet{uchida68} with his coronal fast-mode MHD wave interpretation of Moreton waves. Though this interpretation has been confirmed for several case studies \\citep[e.g.][]{thompson99, warmuth01, pohjolainen01, vrsnak02, khan02, vrsnak06, veronig06, muhr10}, the statistical characteristics of EUV and Moreton waves are quite different. In particular, EIT waves generally propagate at much lower velocities, $v\\approx200$--400~km\\,s$^{-1}$, than Moreton waves, $v\\approx1000$~km\\,s$^{-1}$ \\citep[e.g.][]{klassen00, biesecker02, thompson09}. Thus, there is still an ongoing debate of whether EIT waves are really the coronal counterparts of Moreton waves. Additionally, it is still an open discussion whether they are caused by the flare explosive energy release or the erupting CME \\citep[e.g.][]{warmuth01, warmuth04b, zhukov04, cliver05, vrsnak08} and whether they are waves at all or rather propagating disturbances related to magnetic field line opening and restructuring associated with the CME liftoff \\citep[e.g.][]{delanee99, chen02, attrill07, wills07}.  Numerical simulations of EIT waves as fast-mode magnetohydrodynamical (MHD) waves resulted in wave phenomena mimicking observational data quite closely \\citep{wang00, wu01, ofman02, ofman07}. An interesting by-product of the simulations by \\cite{ofman02} and \\cite{terradas04} are stationary brightenings in the vicinity of the launch sites, in line with observations of such areas at the fronts of EIT waves \\citep{delanee99, delanee00, attrill07}. \\citet{delannee07} suggested them to be a result of magnetic field restructuring during the CME liftoff generated either by Joule heating or an increase in density due to plasma compression. Simulations by \\citet{cohen09} favored the latter interpretation with some additional effect due to increased temperature resulting from plasma compression.  Until the launch of the STEREO \\citep[Solar-Terrestrial Relations Observatory,][]{kaiser08} mission in 2006, with the Extreme UltraViolet Imager \\citep[EUVI;][]{howard08} onboard, observations of EIT waves were drastically limited by a $\\approx12$~min cadence of the EIT instrument in the 195~$\\hbox{\\AA}$ passband. The identically built instruments EUVI-A and EUVI-B onboard the twin STEREO spacecraft observe the solar corona in four different EUV passbands with a high observing cadence (up to 75~s) and a large field-of-view (up to 1.7~R$_{\\hbox{$\\odot$}}$) from two different vantage points.  Studies of large-scale waves with STEREO/EUVI reveal velocities in the range of $\\approx200$--400~km\\,s$^{-1}$ and a decelerating character consistent with freely propagating large-amplitude MHD fast-mode waves \\citep{veronig08, veronig10, long08, kienreich09}. With simultaneous observations by the two STEREO spacecraft, it was for the first time possible to analyze the 3D nature of EIT waves with stereoscopic techniques \\citep{kienreich09, patsourakos09, patsourakos09b, ma09, temmer11}. \\citet{veronig08}, \\citet{long08} and \\citet{gopalswamy09} reported the refraction and reflection of a wave at the border of a coronal hole providing strong evidence for the wave nature of the phenomenon. However, \\citet{attrill10} questioned these findings, favoring instead the hypothesis of two simultaneously launched waves from two distinctively separated sites. In the EUVI wave event of 2010 January 17, studied by \\citet{veronig10}, first observations of the full 3D wave dome and its lateral and radial expansion could be performed. In a recent study that became available during the revision process of the present paper, \\citet{long11} analyze the kinematical aspects and the evolution of the full width at half maximum (FWHM) and integral of the 2007 May 19 and 2009 February 13 events with similar methods. Therefore we will discuss and compare their findings with our results in section~\\ref{sec:discussion}. For recent reviews regarding the kinematics of large-scale EUV waves, their morphology and relationship to associated solar phenomena we refer to \\citet{wills10} and \\citet{gallagher10}.  In this paper, we study high-cadence EUV observations of four well defined EUV waves observed by STEREO/EUVI that occurred between May 2007 and April 2010. These waves are special due to their pronounced amplitudes although they occurred during the extreme current solar minimum. Our study has the following main aims: \\begin{enumerate} \\item We apply and compare two different methods to derive the wave kinematics: visual tracking of the foremost part of wave front and a method based on the intensity profiles of the propagating perturbation. \\item Based on the perturbation profiles, we derive quantities which provide insight into the physical character of the phenomenon in terms of Mach numbers as well as evolution of the amplitude, width and integrated intensity of the disturbances. \\item The perturbation profiles are further studied with respect to associated phenomena, such as stationary brightenings, coronal dimmings and rarefaction regions behind the wave fronts. \\end{enumerate} ",
        "conclusions": "\\label{sec:discussion} \\begin{enumerate} \\item We analyzed four well pronounced EUV wave events observed by the STEREO EUVI telescopes in order to compare two different kinematical analysis techniques, the generally used visual tracking method and the semi-automated perturbation profiles. The differences in the determined positions of the leading edge of the wave fronts using both methods are maximal 40~Mm. Thus, we conclude that the perturbation profiles are a suitable method to analyze large-scale waves. The big advantage of the perturbation profile method is the higher degree of automatization and reproductiveness due to the objective measurements of the wave location. Nevertheless, there are some restrictions. In the later evolution phase, the wave fronts become more irregular showing changes in shapes and propagation direction. These effects are not considered in the profile method. Thus, due to the lower amplitude and the smearing of irregular fronts the wave profile is no longer distinct against the background. These properties lead to a systematic underestimation of the distance of the wave compared to the visual tracking method. Additionally, the diffuse bright fronts are easier to identify in RR images due to the fact that they have a higher contrast than BR images.  \\item The determined propagation velocity derived from both methods match within the error limits. For the four wave events under study, we obtain mean velocity values in the range of $220-350$~km\\,s$^{-1}$, i.e. within the velocity range for fast magnetosonic waves during quiet Sun conditions \\citep[e.g.][]{mann99}. Three of the four kinematical curves show only small deceleration values of $-6$ up to $-13$~m~s$^{-2}$, which lie within the error range. For the fastest event of 2007 May 19, the obtained deceleration of $-193$~m~s$^{-2}$ is significant.  \\citet{long11} has analyzed the kinematics of the May 19, 2007 and February 13, 2009 events with a similar perturbation profile method. They find for the May 19, 2007 event the following values: start velocity $v_0$=447$\\pm$87~km~s$^{-1}$, acceleration $a$=--256$\\pm$134~m~s$^{-2}$, a broadening of the pulse width from 50 to 200~Mm and a decrease of the peak amplitude intensity from 60 to 15\\% of the background values. Our findings are $v_0$=429$\\pm$58~km~s$^{-1}$, $a$=--85$\\pm$60~m~s$^{-2}$, and similar results for the broadening of the pulse and their intensity evolution. For the second event of February 13, 2009, \\citet{long11} find $v_0$=274$\\pm$53~km~s$^{-1}$, $a$=--49$\\pm$34~m~s$^{-2}$ while our results are $v_0$=261$\\pm$72~km~s$^{-1}$, $a$=--24$\\pm$30~m~s$^{-2}$. The profile broadening and the intensity evolution is again similar. For the February 13, 2009 event the outcomes of both studies are consistent within the error limits. The results of May 19, 2007 show differences in the acceleration values which can be explained by taking into account that the method used by \\citet{long11} and ours is not identical. 1) While they are extracting the position of maximum intensity we are using the leading edge of the wave fronts for our analysis. The difference in the acceleration values is thus due to this usage of different features for the position measurements. The maximum intensity peaks are propagating at a slower speed compared to the leading edges, an effect that is expected for a feature that is broadening during its evolution. Due to that the deceleration values are smaller for the leading edge measurements while the starting velocity values are not affected. 2) The overall sector width used in \\citet{long11} is varying from event to event while we are using a constant angular width of 45$\\degr$ in each event. 3) The main direction of the calculated sector as well as the the initiation centers are not exactly the same (the uncertainty is about 20~Mm).  \\item Considering the events under study as low-amplitude MHD fast-mode waves the Mach numbers and wave velocity values are proportional to each other, $M_{ms}=v/v_{ms}$. Thus, from the derived Mach number evolution (from the maximum value down to $M_{ms}\\approx 1$) we expect a decrease of the propagation velocity by $\\Delta v \\approx100$ km\\,s$^{-1}$ (2007 May 19), $\\approx35$ km\\,s$^{-1}$ (2009 February 13), $\\approx60$ km\\,s$^{-1}$ (2010 January 17) and $\\approx30$ km\\,s$^{-1}$ (2010 April 29). The velocity changes we derive from the least-square quadratic fits reveal $\\approx 200$ km\\,s$^{-1}$ (2007 May 19), $\\approx40$ km\\,s$^{-1}$ (2009 February 13), $\\approx20$ km\\,s$^{-1}$ (2010 January 17) and $\\approx30$ km\\,s$^{-1}$ (2010 April 29). For the 2009 February 13 and 2010 April 29 events, the derived velocity changes are in the same order of $\\approx30$ km\\,s$^{-1}$ as the error on the velocity determination. Hence, the weak deceleration is hidden in the measurement uncertainties. For the 2007 May 19 event the result of $200$ km\\,s$^{-1}$ clearly exceeds the error in velocity of $\\approx70$ km\\,s$^{-1}$, thus the deceleration is observable and significant. According to the Mach number evolution for the 2010 January 17 event, an observable deceleration of the propagating wave is expected. The determined velocity change of $\\approx60$ km\\,s$^{-1}$ exceeds the velocity error of $20$ km\\,s$^{-1}$, thus masking of the deceleration by data scatter is unlikely. Nevertheless, from the kinematical measurements obtained with the visual tracking and the perturbation profiles we do not derive a significant deceleration, as it would be expected from the peak Mach number estimated for this event.   \\item Each event can be followed over a period of at least 30 minutes, during which we first observe an intensification of the wave, which is followed by a steady decrease and broadening of the wave pulses as extracted from the perturbation profiles. From this we derived useful information on the characteristics of the wave pulses, like the peak amplitude, the width, and the integral of the wave pulse. The evolution of the pulse width shows a clear broadening over time by a factor of 2--3, similar to the broadening of the FWHM values. These values are typical for large-scale waves \\citep{warmuth04b, wills06, warmuth10, veronig10}. The integral below the perturbation profile basically remains constant (2007 May 19 and 2010 April 29) or decreases by a factor of 3\\,--\\,4 (2009 February 13 and 2010 January 17). We want to stress that the combination of profile broadening and amplitude decay leading to a constant integral below the disturbance profile is consistent with the characteristics of a freely propagating wave \\citep{landau87}.  \\item The peak magnetosonic Mach numbers derived for the events under study are in the range of $M_{\\rm{ms}}=1.08$\\,--\\,1.21, indicative of the evolution of weak coronal shocks. The magnetosonic Mach number for 2007 May 19 was determined via two different approaches leading to similar results in a range of $M_{\\rm{ms}}=1.2$\\,--\\,1.4. Thus, the usage of the perturbation maximum amplitude for the determination of the magnetosonic Mach number seems reasonable. \\citet{grechnev11} discussed the problem of projection effects in measurements of EUV wave. Since we do not observe the perturbation from a view-point that is located directly above the propagating disturbance but from an inclined position, the intensity values are always underestimated. This is probably also the reason why we do observe Gaussian wave pulse profiles instead of a sharp edge, expected in the case of a shocked disturbance.  \\item We observe three different distinct features in the perturbation profiles occurring after the wave passage. In three of four events deep coronal dimmings are prominent at small distances from the eruption center. Only for the 2010, January 17 event it is hardly present in the perturbation profiles of the used sector. However, an inspection of Figure~\\ref{img:20100117195Bbaseratio_rot0_overview} reveals the existence of coronal dimmings in the near vicinity of the initiation site most prominent to the western sectors close to the limb and also off-limb.  The second distinct feature in the wake of three of the four the EUV waves are stationary brightenings. These stationary parts appear after the waves' passage at distances of 100--300~Mm and fade only slowly away. Due to the fact that the the most intense wave profile always forms one frame before the stationary brightening appears, these parts seem to be an obstacle for the wave front. These stationary brightenings are not circumferential around the eruption center but confined to specific sectors, which is best observed for the event of 2009 February 13. \\citet{cohen09} did numerical simulations on this event and showed that stationary brightenings are developing at the outer edges of the core coronal dimmings. Observations by \\citet{delanee99, delanee00} and \\citet{delannee08} as well as simulations by \\citet{chen05b} indicate that these stationary brightenings appear at the locations where the connectivity of the magnetic field lines changes. Since these stationary brightenings are located in front of the deep core dimming regions (interpreted as the footprints of the expanding CME), are restricted in their expansion (their maximal distance is given by the first wave fronts location) and appear only in confined sectors after the waves' passage, they can be interpreted as a signature of the CME expanding flanks.   The third distinct feature are rarefaction regions observable in two of four wave events, on 2007 May 19 and 2010 April 29. The rarefaction region behind the wave pulse is a prospective feature that develops trailing a large-amplitude perturbation \\citep{landau87}. Indeed, simulations by \\citet{cohen09} revealed regions, which are dim in intensity and propagate after the bright wave front. \\end{enumerate}  The EUV wave velocities are in the range of fast magnetosonic waves in the quiet solar corona and show either a decelerating characteristics or constant velocity (within the measurement uncertainties) during propagation. The perturbation profiles reveal intensification, broadening and decay of the propagating wave pulse, while the integral evolution is either constant, or decreasing. The small Mach numbers are indicative of a weak coronal shock. Stationary brightenings formed in confined sectors at the outer edge of the core coronal dimming can be interpreted as signatures of the expanding CME flanks. The first observable wave fronts are always located right in front of them. Thus, we conclude that the EUV wave events can be interpreted as freely propagating, weak shock fast-mode MHD waves initiated by the CME lateral expansion."
    },
    "1107/1107.2844.txt": {
        "abstract": "We  reanalyze the problem of Li abundances in red giants of nearly solar metallicity. After an outline of the problems affecting our knowledge of the Li content in low-mass stars ($M \\le 3 M_{\\odot}$), we discuss deep-mixing models for the RGB stages suitable to account for the observed trends and for the correlated variations of the carbon isotope ratio; we find that Li destruction in these phases is limited to masses below about 2.3 $M_{\\odot}$. Subsequently, we concentrate on the final stages of evolution for both O-rich and C-rich AGB stars. Here, the constraints on extra-mixing phenomena previously derived from heavier nuclei (from C to Al), coupled to recent updates in stellar structure models (including both the input physics and the set of reaction rates used), are suitable to account for the observations of Li abundances below $A$(Li) $\\equiv~ \\log \\epsilon$(Li) $\\simeq$ 1.5 (and sometimes more). Also their relations with other nucleosynthesis signatures of AGB phases (like the  abundance of F, the C/O and $^{12}$C/$^{13}$C ratios) can be explained. This requires generally moderate efficiencies ($\\dot M \\lesssim 0.3 - 0.5  \\times 10^{-6} M_{\\odot}$/yr) for non-convective mass transport. At such rates, slow extra-mixing does not modify remarkably Li abundances in early-AGB phases; on the other hand, faster mixing encounters a physical limit in destroying Li, set by the mixing velocity. Beyond this limit, Li starts to be produced; therefore its destruction on the AGB is modest. Li is then significantly produced by the third dredge up. We also show that effective circulation episodes, while not destroying Li, would easily bring the $^{12}$C/$^{13}$C ratios to equilibrium, contrary to the evidence in most AGB stars, and would burn F beyond the limits shown by C(N) giants. Hence, we do not confirm the common idea that efficient extra-mixing drastically reduces the Li content of C-stars with respect to K-M giants. This misleading appearance is induced by biases in the data, namely: i) the difficulty of measuring very low Li abundances in O-rich AGB stars, due to the presence of TiO bands; and ii) the fact that many, relatively massive ($M >3 M_{\\odot}$) K- and M-type giants may remain Li rich, not evolving to the C-rich stages. Efficient extra-mixing on the AGB is instead typical of very low masses ($M \\lesssim$ 1.5$M_{\\odot}$). It also characterizes CJ stars, where it produces Li and reduces F and the carbon isotope ratio, as observed in these peculiar objects.  {\\noindent {\\bf key words}: stars: abundances -- stars: evolution of -- nucleosynthesis -- stars: RGB and AGB -- light elements.} ",
        "introduction": "Li abundances in solar metallicity low mass stars (hereafter LMS) are distinctively peculiar and their spread reveals the inadequacy of standard stellar codes, not including rotation and assumptions for non-convective forms of mixing. $^7$Li is produced from electron captures on the short-lived $^7$Be ($t_{1/2}$ = 53 days in laboratory) and, due to its fragility, is sensitive to any mixing event connecting the envelope  with warm layers ($T \\ge 2.5 - 3 \\times 10^6$ K), where it easily undergoes p-captures; its production and destruction are also critically dependent on the H-burning rates controlling the temperature stratification of a star.  A large number of observational and theoretical works tried to interpret the observations of Li on the Main Sequence (hereafter MS) and on the Red Giant Branch (RGB), for both Population II and Population I stars. By contrast, the number of studies addressing its behavior in the stages of the Asymptotic Giant Branch (AGB), characterized by the presence of thermal instabilities, or {\\it pulses} (TPs) from the He shell, is smaller and the observational data are less numerous. Nevertheless, these stages play a crucial role because they show us the integral of all the complex phenomena affecting the Li concentration throughout the evolution of a star. They are, moreover, the site where fresh nucleosynthesis occurs in He-rich layers coupled to convective mixing episodes, the so-called third dredge-up (TDU), increasing the envelope abundance of carbon (up to C/O $>$ 1 for C stars), of fluorine, of technetium and of other neutron-capture isotopes; these independently-produced nuclei can provide tools to understand what is happening to Li. In this paper we plan therefore to reanalyze the abundances of Li in evolved stars, with emphasis on the AGB stages.  Non-convective transport of H-burning products was suggested by many authors to be required for explaining the Li abundances and other chemical anomalies in evolved stars \\citep[see e.g.][]{was95,sb99,cdn98,b+07,b+10,clan10}. Various physical processes have been explored for driving such a transport: reviews can be found in \\citet{pinson97} and \\citet{talon08}. Special attention was dedicated, in recent years, to thermohaline diffusion and rotational shear \\citep[see e.g.][]{chala10}.  %Whatever mechanism causes mixing, %it is expected to account also for other embarrassments of present stellar evolution, %like e.g. the angular momentum transport necessary to explain the present slow %rotational rate of the solar core \\citep{t95,dp06}.  The mixing phenomena induced directly by rotation were shown by \\citet{pala6} to be rather ineffective in evolved stars and the recent analysis by \\citet{chala10} specifies that their role is mainly that of accounting for star-to-star scatters in Li abundances and of favoring an early contamination of the envelope with the materials brought up by diffusion mechanisms. Rotational mixing can be instead important on the MS, where its low efficiency is compensated by the long available time scales \\citep{clan10}.  On the other hand, thermohaline diffusion, as discussed in \\citet{egg1,egg2} and in \\citet{cz07}, has the advantage that it descends from an understood mechanism, namely $^3$He burning into $^4$He and two protons, reducing the molecular weight and hence inducing mass readjustments. For this reason, it has rapidly become the most commonly assumed process of transport. However, recently severe doubts have been advanced also on its effectiveness \\citep{den10,dm10}; the inferred velocities of the transport (and/or the aspect ratios of the salty \"fingers\" found in hydrodynamical simulations) were found to be too small by large factors to yield significant changes in the envelope abundances. This subject must be clearly examined with different hydrodynamical codes by others before final conclusions can be drawn. Nevertheless we note that a recent study of the globular cluster M3 confirms that high mixing velocities and large aspect ratios of the mixing structures, beyond the reach of thermohaline diffusion, are required \\citep{angelou}.  %An alternative might be offered, in evolved stars, by magnetically-induced transport. %Revising works by \\citet{tay,ache,pt} and others, \\citet{sp1} %and \\citet{sp2} proposed that angular momentum transport in radiative layers might %be provided by a special dynamo-like mechanism (which has been since called Tayler-Spruit %dynamo). Its effects on mixing and angular momentum transport were analyzed, e.g., by %\\citet{mm4} and \\citet{egg}. They have been subsequently revised by \\citet{dp06}, %who pointed out that, although this scenario is important in stellar %evolution, other mechanisms must be active since the Main Sequence to %explain the spin-down of stellar cores.  In the above uncertain situation, a few recent papers suggested that magnetic fields might play a role in stellar mixing through the buoyancy of light magnetized bubbles, either alone \\citep{b+07,nor08} or in synergy with thermohaline diffusion \\citep{den09,dm10}. Qualitative arguments in favor of this scenario have been outlined by \\citet{b+10} and by \\citet{p+11}, hereafter Paper 1. In any case, the above discussion makes clear that attributing the abundance anomalies of evolved stars to a specific physical mechanism might still be a long shot. Hence we prefer to adopt a parametric model similar to the ones used in \\citet{nol03} and in Paper 1, in an attempt to verify if mixing episodes of whatever nature (characterized by the parameters derived in Paper 1 from an explanation of CNO isotopic admixtures in evolved stars) can also account for the behavior of Li.  The present work is organized as follows. In Section 2 we recall the general history of the Li evolution in LMS  before the AGB phase. We also present a set of relevant observational data derived from the current literature, which can be used to constrain the nucleosynthesis calculations. In this way we show that deep-mixing schemes based on a range of (moderate) circulation rates can easily account for most of them. In Section 3 we discuss observations on both O-rich and C-rich TP-AGB stars. In Section 4 we address some issues related to the time scales for mixing and for $^7$Be decay, which are relevant for Li production and destruction. Then Sections 5 and 6 afford the interpretation of Li abundances and of their relations with other observed nuclei in O-rich and C-rich AGB stars, respectively. For this task we use the stellar models and the extra-mixing scheme already discussed in Paper 1. Some general conclusions are then drawn in Section 7. We also add, in Appendix I, a brief outline of our knowledge on the most relevant reaction rates for our analysis. ",
        "conclusions": "In this paper we have analyzed those evolved red giants that generally show Li abundances $A$(Li) $\\lesssim$ 1.5 in their spectra, i.e. envelope concentrations lower than (or equal to) the typical maximum values provided by stellar models after the FDU. We attributed the Li consumption to the occurrence of extra-mixing mechanisms, which we parameterized, in line with what was done in Paper 1, on the basis of the efficiency of mass transport and of the maximum temperature of the mixed layers. At solar metallicity, circulation rates from low to moderate ($\\dot M_6 = 0.03 - 0.3$) are sufficient to account for the observations of Li and $^{12}$C/$^{13}$C ratios in RGB stars. This would roughly correspond, in a diffusive treatment, to diffusion coefficients $D_{mix}$ of the order of 3$\\times10^6$ to 3$\\times$10$^7$ cm$^2$/sec, i.e. on average a factor-of-ten smaller than found by \\citet{dw} for Population II stars. This might be a confirmation of the known decrease of deep mixing efficiency with increasing metallicity.  The reproduction of the Li observations in red giants requires consideration of a rather large spread of Li abundances at FDU, as can be found in the current literature. These abundances, in their turn, imply that non-convective mixing has previously operated on the MS at different efficiencies, probably due to rotational effects, thus dispersing the Li abundance that is induced in the envelope by FDU, which would be otherwise confined in a small range ($A$(Li) = 1 $-$ 1.5).  The occurrence of extra-mixing in RGB phases is possible after the advancement of the H-burning shell has erased the chemical discontinuity created by FDU. With the reaction rates revisions discussed in Paper 1 we find that only in stars below about 2.3$M_{\\odot}$ this can be obtained before the RGB tip. More massive stars, therefore, cannot experience deep-mixing episodes during their ascent to the RGB. On the other hand, during core He-burning and through the major part of the ascent to the early-AGB, the envelope remains rather far from the H-burning shell; any extra-mixing phenomenon is therefore bound to restart only later, when the star is again close to or directly on the Hayashi track. This includes the last $\\simeq$ 1 Myr of the early-AGB and the subsequent TP-AGB. Globally, the  available time (about 4 Myr) is typically a factor-of-ten shorter than for the RGB, so that if extra-mixing occurs at the same small or moderate rates, little Li is destroyed. This fact is in agreement with what was shown by \\citet{guan09}, although in their analysis the subsequent Li production from TDU is prevented by the very low abundance of $^3$He they considered. Conversely, when convective TDU episodes start, they reproduce Li, rapidly saving to the envelope whatever $^7$Be has remained above the H shell. On the other hand, we showed that any faster extra-mixing episode (at $\\dot M_6 \\gtrsim$ 1) does not destroy more Li, but actually re-produces it, because its overturn time scale becomes, at least in the external layers of the radiative zone, faster than that for Be decay.  We therefore conclude that effective Li destruction on the AGB, suitable to establish an evolutionary link between Li-rich K or M giants and Li-poor C(N) stars is not physically possible. The common interpretation that a considerable Li destruction must occur in between these evolutionary stages is wrong. On the contrary, stars evolve along the AGB, before the TPs, at almost constant Li. Li-poor C(N) giants must therefore descend from Li-poor early-AGB stars. These last are observed in very limited numbers only because of the presence of TiO bands, preventing the measurements of low Li abundances, at least in cool M giants. Conversely, Li-rich M giants (statistically much more abundant that Li-rich C(N) stars) should owe their high Li content to the fact that their mass is high enough to prevent extra-mixing to occur on the RGB ($M > 2.3 M_{\\odot}$). Only the less massive among them will become C-rich, the others will remain O-rich (due to their too large envelope masses, for which TDU cannot induce the C/O $>$ 1 condition). Hence, two different phenomena conjure in creating strong selection effects in the comparisons between O-rich and C-rich AGB stars, inducing the false appearance of a reduction of Li along the AGB.  The above scenario is confirmed by the limited data available for Tc (in the sense that, with the moderate extra-mixing efficiencies mentioned, all Tc-rich AGB stars are fitted by models that are crossing the TP-AGB stage, where Tc is mixed to the envelope by TDU).  A more quantitative check is possible thanks to F abundances. Extra-mixing at rapid rates and reaching down to hot temperatures destroys F, while most C(N) stars appear to be F-rich. Hence, the trend of fluorine confirms that only moderate transport rates (leaving unchanged the abundances of F dredged up by TDU) must characterize the majority of C(N) stars. In this scenario, CJ stars seem to be those few objects that experienced fast mixing: its occurrence indeed produces Li, and they are Li-rich; it destroys $^{12}$C producing $^{13}$C, and they have low $^{12}$C/$^{13}$C ratios. It also destroys F, and CJ stars are known to be F-poor.  \\begin{figure}[h!!] \\centerline{{\\includegraphics[height=7.8cm, width=7.5cm]{lifig13a.eps}} {\\includegraphics[height=7.8cm, width=7.5cm]{lifig13b.eps}}} \\caption{The effects of a reduction by 30\\% of the rate for Be decay (bold continuous curve), as compared to the presently accepted value (normal continuous curve). The allowed variation accounts for the uncertainties mentioned for this rate in Appendix I. As is clear from the plot, the balance between Li destruction and production, due to the competition of the time scale for mixing with that for decay, is changed only partially on the AGB (right panel) and even less on the RGB (left panel).} \\label{lifig13} \\end{figure}  All the above conclusions are obviously tentative, as the interpretation of the data is not straightforward and the statistics is not large. Cautions are also due to the uncertainties affecting the rate for $^7$Be decay, as discussed in the Appendix. One has to notice that the H-burning shell in evolved stages has a higher temperature and a lower density than the central regions of the present Sun. Hence the plasma conditions should make Be decay through the capture of free electrons even more difficult than in the Sun. Obviously, a longer Be lifetime would require changes in our discussion about the competition between the time scale for decay and that for mixing overturn. However, if very special effects in the plasma (see below) can be neglected, then the effects of a possibly longer half-life for $^7$Be are not expected to be dramatic. In order to give a more quantitative support to this statement, Figure 13 (bold continuous line) shows the effects of an increase by 30\\% of this half life, according to the suggestions by \\citet{shsh}; the presently accepted value is shown by a normal continuous line. Figure 13 refers to the most critical case, that of a 2$M_{\\odot}$ star. It shows that the change does  modify drastically the competition with the overturn time scale for mixing on the RGB; on the AGB it is more relevant, implying an increase of the effective area for Li production (where $\\tau_{mix} \\lesssim \\tau_{e^-}$), which is however not particularly crucial, given all the other uncertainties present. The main source of difficulty in modeling the Li abundances in red giants seems therefore to be linked to the astrophysical scenario, not to the nuclear input data. A different situation might occur if the presence of magnetic fields can modify the energy distribution of the particles inside the plasma, inducing Lorentz forces.  However, no nuclear physics calculation seems to have considered this possibility so far, so that the rate for Be decay in such a situation remains essentially unknown. A more complete analysis of this issue will be performed in the future paper dedicated to Li-rich stars.  \\newpage  {\\bf APPENDIX I Nuclear Processes Controlling the Nucleosynthesis of Li}  Most nuclear reaction rates used in the present work are based on a recently published review of solar fusion cross sections \\citep{ade11}. For some reactions not included there, the recommendations by NACRE \\citep{angulo} are adopted.  We report here a brief analysis of the present situation, underlining  the uncertainties affecting nuclear data that might have an impact on the synthesis of lithium in stars.  {\\bf 1. $^2$H(p,$\\gamma$)$^3$He and $^3$He($^3$He,2p)$^4$He.} The recommendations by \\citet{ade11} for the astrophysical S-factors of these reactions are S$_{12}$(0) = 0.214$^{+0.017}_{-0.016}$ eV b and S$_{33}$(0) = 5.21$\\pm$0.27 MeV b, respectively. These reactions have been well constrained experimentally down to the Gamow peak in the Sun \\citep{case,bone} and there appears to be no room for anything different than an almost instantaneous conversion of deuterium into $^3$He and the following production of a rather well defined abundance of $^4$He.  {\\bf 2. $^3$He($^4$He,$\\gamma$)$^7$Be.} The recommendation by \\citet{ade11} is S34(0)=0.56$\\pm$0.02(exp.)$\\pm$0.02(theor.) keV b, where the more recent measurements by \\citet{singh}, \\citet{br07}, \\citet{costa}, and \\citet{dileva} are considered. A closer look at the data shows a  slight tension between them, possibly indicating some discrepancy due to systematic uncertainties. Recently, a new {\\it ab-initio} calculation \\citep{neff} has been presented, where a realistic nucleon-nucleon interaction is used to calculate the cross section. The resulting S$_{34}$ is 0.593 keV b. The agreement with the data by \\citet{br07} and \\citet{dileva} is very good, both regarding the absolute value and the energy dependence of the S factor; instead, the agreement with the results from LUNA \\citep{costa} and from the Weizman Institute \\citep{singh} is at a 2$\\sigma$ level or worse. An extrapolation leading to S$_{34}$=0.59 keV b should therefore be considered as a realistic alternative to the quoted recommended value (this would however imply only an increase by 5\\%). In the paper we adopt, for a question of uniformity, the \\citet{ade11} choice; the possible increase is too low to have any significant effect.  {\\bf 3. $^7$Be(e$^-$ $\\nu$)$^7$Li.} Here the recommendation by \\citet{ade11} is essentially based on the work made by Bahcall and collaborators \\citep{bahc69}. In turn, their calculations were based on the paper by \\citet{iben}, where a partial ionization of $^7$Be in the Sun was predicted. Therefore, the rate now adopted includes both a bound state and a continuum contribution to the decay, with a total-to-continuum capture ratio of 1.217. A rather different result was obtained by \\citet{shsh}, who predicted a fully ionized condition for $^7$Be in the Sun. The missing contribution from bound electrons to the capture rate was partially compensated by a larger contribution from the continuum. Summing up, an increase of the lifetime by only 20 to 30\\% resulted, as compared to the above-mentioned recommendations. The situation is now even more complicated, as \\citet{qs9} recently found a $^7$Be lifetime shorter by about 10\\%, due to the use of a modified Debye-Hueckel screening potential. On the contrary, \\citet{bahc} excluded the validity of a modified Debye-Hueckel potential for a dense stellar plasma. In conclusion, the real questions in this case are: i) which is the degree of ionization of $^7$Be in a stellar plasma, e.g. in the Sun core? ii) Further, as the $T$ conditions above the H-burning shell in evolved stars span a large range (from a few $\\times$10$^7$ to a few $\\times$10$^6$ K) and as the density is lower than in the Sun, how does the $e^-$-capture rate react? (H-shell conditions, in particular, might modify considerably the equilibrium between bound and free electrons, and even the number of free electrons contributing to the capture). The large uncertainties, both of observational and theoretical nature, affecting Li abundances in stars have also to include this poor understanding of the basic nuclear input data.  {\\bf 4. $^7$Be(p,$\\gamma$)$^8$B.} In this case the recommended value by \\citet{ade11} is $S_{17}$(0) = 20.8$\\pm$7(exp.)$\\pm$1.4(theor.) eV b. This value is estimated considering only direct measurements (whose results show some mutual discrepancy); indirect (Coulomb dissociation) measurements were not taken into account. The comparison between direct and indirect measurements is debated \\citep{esbe,gai06}, but the overall impact on the $^7$Li abundance should be negligible, when compared to the uncertainty deriving from the reactions $^3$He+$^4$He and $^7$Be+e$^-$.  {\\bf 5. $^7$Li(p,$\\gamma$)$^8$Be.} This reaction has not been included in the compilation by \\citet{ade11}, so we use the NACRE recommendation \\citep{angulo}, which is $S_{71\\gamma}$(0)= 0.3 keV b. The measurements by \\citet{spraker} later suggested a doubling of this rate; however its impact on Li destruction is minimal, due to the prevailing role of the (p,$\\alpha$) channel.  {\\bf 6.  $^7$Li(p,$\\alpha$)$^4$He.} Also for this reaction, we use the NACRE value ($S_{71\\alpha}$(0)= 55.6$^{+0.8}_{-1.7}$ keV b, later confirmed by \\citet{cruz}. One should note, however, that an anomalous screening is observed in this case, whose origin is not well understood. This might have an impact on the stellar rate.  {\\bf Acknowledgements} We are indebted to S. Randich and L. Magrini for many useful discussions on spectroscopic data for Li abundances in open clusters and in the Sun. We then owe a lot to OUR longstanding collaborations with R. Gallino, K.M. Nollet, L. Piersanti, O. Straniero and G.J. Wasserburg. S.P and M.B. are grateful to Group III of INFN for support. C.A. and S.C. were partially supported by the Spanish grants AYA2008-04211-C02-02 and FPA2008-03908 from the MEC. S.U. acknowledges support from the Austrian Science Fund (FWF) under project number P~22911-N16."
    },
    "1107/1107.4571_arXiv.txt": {
        "abstract": " ",
        "introduction": "} The Initial Mass Function (IMF) describes the number of stars formed in a stellar system as a function of stellar mass, and is a fundamental property of all stellar populations.  As a statistical measure of the end result of the star formation process, the IMF is a key observable for testing theoretical models of star formation.  As a succinct characterization of the fundamental components of a stellar population, the MF also serves to inform our understanding of the structure and dynamical evolution of stellar clusters, the Milky Way and other galaxies.  Numerous physical processes have been identified which may influence the shape of the IMF, such as: gravitational fragmentation of collapsing molecular cores \\citep{Klessen1998}; competitive accretion between multiple stars inhabiting the same mass reservoir \\citep{Larson1992}; the truncation of mass accretion due to radiative or dynamical feedback \\citep{Silk1995}; dynamical interactions between stars in a clustered environment \\citep{Reipurth2001}; and the production of a clump mass spectrum by turbulent flows within molecular clouds \\citep{Padoan2002}. The efficiency of each mechanism could also depend on other physical variables, such as the metallicity and magnetic field strength of the parent molecular cloud, the local stellar density, or the intensity of the surrounding radiation field.  These effects may ultimately result in observable MF variations as a function of environment.  The effort to provide observational constraints on the form of the IMF can be traced back to the `Luminosity Curve' measured by \\citet{Kapteyn1914}, which determined the relative numbers of B type stars as a function of absolute magnitude.  \\citet{Salpeter1955} subsequently produced a measurement of the IMF for high mass stars which has remained essentially unchanged to the present day.  \\citet{Salpeter1955} found that the shape of the IMF took the form of a power law, which can be expressed as:  \\begin{equation} \\Phi(logm) = dN/d\\log m \\propto m^{-\\Gamma}, \\label{eqn:salpeter} \\end{equation}  \\noindent where $m$ is the mass of a star and N is the number of stars in some logarithmic mass range $log m+d log m$.   \\citet{Salpeter1955} inferred a value of $\\Gamma=$1.35, which has come to be known as the `Salpeter slope'.  While numerous observational studies have found the Salpeter slope to be a good description of the IMF in the super-solar mass regime, one of the first measurements of the low-mass IMF revealed that solar-mass and sub-solar mass stars are slightly less numerous than might be expected from an extrapolation of the Salpeter slope \\citep{Miller1979}.  Changes in the slope of the IMF can be expressed within the power-law formalism by allowing different mass regimes to possess distinct power-law slopes, as in the seminal `broken power-law' IMF derived by \\citet{Kroupa1993}.   \\citet{Miller1979} adopted a different approach, describing the IMF over a large mass range with a single analytical expression, a gaussian in $log (m)$, often known as a `log-normal' function  \\begin{equation} \\phi(m) \\sim e^{-\\frac{({\\log}~m-{{\\log}~m_c})^2}{2\\sigma^2}} \\label{eqn:lognormal} \\end{equation}  \\noindent where the variable $m_c$ fixes the peak of the IMF (in $log (m)$ space), and $\\sigma$ characterizes the peak's width.  Distinctions are often drawn between the power-law and log-normal characterizations of the IMF, but these differences are currently entirely in the realm of theory, not observation: \\citet{Dabringhausen2008} have shown that the log-normal IMF advanced by \\citet{Chabrier2005a} is extremely similar to a two-part power-law, hence distinguishing between a Kroupa-type broken power-law or Chabrier-type log-normal IMF is virtually impossible.  A great deal of observational work has been devoted to characterizing the IMF in a variety of astrophysical environments, across the full range of stellar masses, and extending into the brown dwarf regime.  In a recent review \\citep{Bastian2010}, we provided a overview of recent empirical measurements of the IMF, and evaluated the evidence for systematic IMF variations.  In this contribution we update that review, focusing on recent (2009-2011) IMF measurements in resolved stellar populations, which the upcoming Gaia mission will characterize in exquisite detail.  Specifically, we review recent measurements of the mass function in the extended solar neighborhood (Section 2), in young star forming regions (Section 3), and Galactic open/globular clusters (Section 4).  We conclude in Section 5 by examining the complimentary roles that Gaia and the Large Synoptic Survey Telescope will play in extending and improving IMF studies of resolved stellar populations. ",
        "conclusions": ""
    },
    "1107/1107.4109_arXiv.txt": {
        "abstract": "We use the Expanded Very Large Array to image radio continuum emission from local luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs) in 1 GHz windows centered at 4.7, 7.2, 29, and 36~GHz.  This allows us to probe the integrated radio spectral energy distribution (SED) of the most energetic galaxies in the local universe. The 4--8~GHz flux densities agree well with previous measurements. They yield spectral indices $\\alpha \\approx -0.67$ (where $F_{\\rm \\nu} \\propto \\nu^{\\alpha}$) with $\\pm 0.15$ (1$\\sigma$) scatter, typical of nonthermal (synchrotron) emission from star-forming galaxies. The contrast of our 4--8~GHz data with literature $1.5$ and $8.4$~GHz flux densities gives further evidence for curvature of the radio SED of U/LIRGs. The SED appears flatter near $\\sim 1$~GHz than near $\\sim 6$~GHz, suggesting significant optical depth effects at the lower frequencies. The high frequency (28--37~GHz) flux densities are low compared to extrapolations from the 4--8~GHz data. We confirm and extend to higher frequency a previously observed deficit of high frequency radio emission for luminous starburst galaxies. ",
        "introduction": "\\label{sec:intro}  Most very luminous galaxies at $z \\sim 0$ emit the bulk of their luminosity in the infrared (IR). These luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs, infrared luminosity $L_{\\rm IR} \\left[ 8-1000\\mu {\\rm m} \\right] > 10^{11}$ and $L_{\\rm IR} > 10^{12}$~$L_{\\odot}$) make up the majority of galaxies with $L_{\\rm bol} \\gtrsim 10^{11.5}$~L$_\\odot$.  These systems often correspond to major mergers and compared to normal star-forming disk galaxies, U/LIRGs exhibit enhanced, concentrated molecular gas content, high star formation efficiency, increased active galactic nuclei (AGN) activity, and disturbed morphologies \\citep{SANDERS96}. Even the closest ULIRGs are $\\gtrsim 100$~Mpc away with compact radio, IR, and millimeter line emission \\citep[e.g.,][]{LONSDALE93,DOWNES98,SAKAMOTO08}, so that high angular resolution (such as that obtained by the EVLA) is required to study their structure.  Nonthermal (synchrotron) radio emission is well-established to track infrared emission \\citep[the \"radio-FIR correlation,\"][]{HELOU85,CONDON92,YUN01} while thermal (bremsstrahlung) radio emission at $\\nu \\gtrsim 4$~GHz gives an unobscured view of ionizing photon production. These emission mechanisms, the indifference of radio waves to dust, and the resolving power of cm-wave interferometers make the radio continuum a natural tool to dissect the detailed structure of U/LIRGs. Deviations from the radio-IR correlation, radio morphology, and brightness can all help distinguish the presence and contribution of AGN to the overall luminosity, though not necessarily uniquely \\citep[][and references in the latter]{CONDON91,LONSDALE95,SAJINA08,LONSDALE06}.  This letter reports first results from our Expanded Very Large Array (EVLA) program to explore the structure of star formation in U/LIRGs using radio continuum emission. We are imaging 22 of the most luminous nearby U/LIRGs ($L_{\\rm IR} > 10^{11.6}$~L$_\\odot$) in the C (4-8~GHz) band and Ka (26.5-40~GHz) bands. Here we report observations from the EVLA's C and D configurations. The EVLA's continuum sensitivity makes it possible to achieve good signal-to-noise ratio, even at high frequencies, with modest time on source. The large instantaneous bandwidth allows us to probe the radio spectral index, $\\alpha$ (defined so that $F_{\\rm \\nu} \\propto \\nu^{\\alpha}$), both within and between bands.  Here explore variations in the radio spectral energy distribution (SED) of our targets and in the radio-IR correlation. In a simple picture dominated by star formation we would expect the C band to be dominated by synchrotron emission with spectral index $\\alpha \\approx -0.8$ while the Ka contains an approximately equal mixture of nonthermal emission and thermal emission, which has $\\nu = -0.1$. For systems dominated by an AGN the spectral index may show wider variation depending on the orientation of the AGN or injection spectrum of electron energies. We close by showing overlays of the Ka-band emission on HST imaging. These highlight the EVLA's improved ability to trace embedded star formation at high resolution. ",
        "conclusions": ""
    },
    "1107/1107.2944_arXiv.txt": {
        "abstract": "Using a dark matter only Constrained Local UniversE Simulation (CLUES) we examine the existence of subhaloes that change their affiliation from one of the two prominent hosts in the Local Group (i.e. the Milky Way and the Andromeda galaxy) to the other, and call these objects \"renegade subhaloes\". In light of recent claims that the two Magellanic Clouds (MCs) may have originated from another region (or even the outskirts) of the Local Group or that they have been spawned by a major merger in the past of the Andromeda galaxy, we investigate the nature of such events. However, we cannot confirm that renegade subhaloes enter as deep into the potential well of their present host nor that they share the most simplest properties with the MCs, namely mass and relative velocity. Our simulation rather suggests that these renegade subhaloes appear to be flying past one host before being pulled into the other. A merger is not required to trigger such an event, it is rather the distinct environment of our simulated Local Group facilitating such behavior. Since just a small fraction of the full $z=0$ subhalo population are renegades, our study indicates that it will be intrinsically difficult to distinguish them despite clear differences in their velocity, radial distribution, shape and spin parameter distributions. ",
        "introduction": "\\label{sec:introduction} In the concordance cosmology, structure forms in a hierarchical, ``bottom-up'' fashion that leads to the accretion of small substructures by large dark matter haloes. The current paradigm holds that substructures orbit within their parent haloes until tidal stripping rips them apart as they sink to the centre of the potential by dynamical friction.   The Local Group of galaxies is an ideal testbed for this kind of near-field cosmology \\citep{Freeman02} as the observational data is better than for any distant galaxy. But as already prompted before, how certain can we be that the Milky Way (MW) or the Andromeda galaxy (M31) are in fact typical galaxies of their mass or luminosity \\citep[cf. ][]{Forero-Romero11}?  These two  galaxies are the dominant members of the Local Group and form a two-body system on a collision course in approximately 2-3 Gyrs \\citep{Cox08,Hoffman07}. Their proximity at the present time indicates that they may already have started to influence each other or their respective satellite populations. Perhaps this influence is related to the (in)famous ``disk of satellites'' \\citep[e.g.][]{Metz08}; there is also the question of how frequently the largest companions  of the MW - the Small and the Large Magellanic Clouds (SMC and LMC) - are found. A recent study by \\citet{Liu10} indicates that only of order 3\\% of Milky Way type galaxies in the SDSS Data Release 7 host two satellites with luminosities similar to the Magellanic Clouds (MC). This result is also supported by \\citet{James10} who investigated 143 luminous spirals in H$\\alpha$ finding that the MW is an unusual galaxy both for the luminosity and the proximity of its two brightest satellites. Therefore, both these satellites are often considered outliers with respect to the full system of satellites orbiting the MW and their formation scenario is a matter of debate \\citep{Besla07, Peebles09, Kallivayalil09, Metz09}.  There are several investigations that suggest that the LMC and SMC are on their first infall as suggested by recent proper motions measurements \\citep{Besla07, Tollerud11} -- as opposed the to recent claims by \\citet{Sales11} that the LMC is not necessarily on its first approach to the MW based upon the Aquarius simulation \\citep{Springel08}; if they are only loosely (if at all) bound satellites of the MW then where did they come from? \\citet{Yang10} deal with the possibility that both the Magellanic Clouds have been expelled from M31 due to a major merger event. The scenario envisioned by them is that they are ejected tidal dwarf galaxies from a previous major merger occurring at the M31 location \\citep{Hammer10}. But there is also the notion that they may simply be accreted objects that formed in the outer reaches of the Local Group \\citep{vandenBergh10, Busha10}. While these studies are still speculative, they nevertheless show that M31 may play a vital role in shaping the orbits of the Magellanic Clouds \\citep[see e.g.][]{Kallivayalil09}. Neither the MW nor M31 may be understood in isolation but only in the context of the (unique) environment of the Local Group and its formation.  In this \\textit{Letter} we use a Constrained Local UniversE Simulation (\\textit{http://www.clues-project.org}) to search for events where subhaloes change their host halo affiliation. These simulations are well suited for this idea as they directly model the Local Group (consisting of the two-body system MW and M31) within the correct environment and in a cosmological framework of the concordance model. We investigate the likelihood that subhaloes were under the influence of one of the two hosts at some previous time while at $z=0$ are within the virial radius of the other host. We colloquially call them ``renegade'' or ``disloyal''  subhaloes. ",
        "conclusions": "\\label{sec:concusions} Motivated by recent claims that the two Magellanic Clouds \\nil{were either formed close to} M31 \\nil{(and ejected from it by} a recent merger) or simply accreted \\nil{from the} outer reaches of the Local Group \\citep{vandenBergh10, Busha10} we \\nil{examined the likelihood of} this situation in Constrained Local UniversE simulations of the Local Group. We find that subhaloes disloyal to their first host exist, and called them renegade subhaloes. \\nil{They are subhaloes accreted by one host halo at a given time, yet at $z=0$ are found within the virial radius of completely different host.}  \\nil{We have examined various physical properties of the renegade population and contrasted them with the full loyal population. We find that significant differences are seen in the relative velocities, sphericity, radial distribution and spin parameter. Furthermore, renegade subhaloes are significantly more anisotropically distributed than the full subhalo population - more renegades are found close to the line connecting our two main galaxies than not.}  Despite the differences \\nil{highlighted above, detecting renegade subhaloes is only possible if they leave an imprint on the full subhalo population - its not enough to have vastly different properties.} One \\nil{thus} has to compare the dashed \\nil{with} the dotted line \\nil{in the figures presented above, to gauge any observable signature. In principle,} an observer will only have access to the set of visible satellite galaxies and hence could generate a distribution akin to the ones presented in \\Fig{fig:PCompare}. But only if loyal satellites show a substantially different behavior to the combined sample will there be a chance to (observationally) confirm their existence \\citep[as, for instance, the difference in the velocity distribution for backsplash and infalling subhaloes reported by][]{Gill05}. We thus conclude that their (observational) detection will be difficult unless orbital information can act as a discriminator.   The \\nil{abundance} of renegade subhaloes \\nil{appears to be} a function of \\nil{both} the distance between the hosts \"sharing\" them (which also corresponds to redshift in our case) and the mass of the hosts. \\nil{Yet despite their existence, it is still a great challenge to explain the origin of the MCs as renegade subhaloes. In our (dark matter only) simulation, renegade subhaloes are not massive enough to realistically be called ``MC''-like. Our inability to find a MC-like renegades, however, may simply be due to the difficulty in producing such massive satellites in the first place. Further, our simulation currently only considers dark matter whereas it has been shown before that baryonic physics will have a certain impact upon the properties of subhaloes \\citep[e.g.][]{Libeskind10,Romano-Diaz08}. Future work that quantifies the frequency of finding MC-like objects in MW sized haloes, and extends this work by then examining the likelihood that these subhaloes become renegades is needed to verify or falsify this formation scenario.}"
    },
    "1107/1107.5049_arXiv.txt": {
        "abstract": "\\ \\\\ The abundance of the most massive objects in the Universe at different epochs is a very sensitive probe of the cosmic background evolution and of the growth history of density perturbations, and could provide a powerful tool to distinguish between a cosmological constant and a dynamical dark energy field. In particular, the recent detection of very massive clusters of galaxies at high redshifts has attracted significant interest as a possible indication of a failure of the standard $\\Lambda $CDM model. Several attempts have been made in order to explain such detections in the context of non-Gaussian scenarios or interacting dark energy models, showing that both these alternative cosmologies predict an enhanced number density of massive clusters at high redshifts, possibly alleviating the tension. However, all the models proposed so far also overpredict the abundance of massive clusters at the present epoch, and are therefore in contrast with observational bounds on the low-redshift halo mass function. In this paper we present for the first time a new class of interacting dark energy models that simultaneously account for an enhanced number density of massive clusters at high redshifts and for both the standard cluster abundance at the present time and the standard power spectrum normalization at CMB. The key feature of this new class of models is the ``bounce\" of the dark energy scalar field on the cosmological constant barrier at relatively recent epochs. We present the background and linear perturbations evolution of the model, showing that the standard amplitude of density perturbations is recovered both at CMB and at the present time, and we demonstrate by means of large N-body simulations that our scenario predicts an enhanced number of massive clusters at high redshifts without affecting the present halo abundance. Such behavior could not arise in non-Gaussian models, and is therefore a characteristic feature of the bouncing coupled dark energy scenario. ",
        "introduction": "\\label{i}  The origin of the observed acceleration of the cosmic expansion \\citep{Riess_etal_1998,Perlmutter_etal_1999,SNLS,Kowalski_etal_2008} represents one of the most relevant open questions in astrophysics and cosmology. Besides the standard $\\Lambda $CDM cosmological model, that assumes a perfectly homogeneous and static energy density represented by a cosmological constant $\\Lambda $ as the source of the acceleration, alternative scenarios based on a dynamical dark energy (DE) field such as {\\em Quintessence} \\citep[][]{Wetterich_1988,Ratra_Peebles_1988, Ferreira_Joyce_1998} and {\\em k-essence} \\citep{ArmendarizPicon_etal_2000, kessence}, or on a modification of General Relativity at large scales \\citep[see \\eg][]{Hu_Sawicki_2007,Appleby_Weller_2010} have been proposed in recent years.  In order to efficiently distinguish among these different possibilities it is therefore crucial to devise observational tests capable of highlighting a possible dynamical nature of the DE field, thereby falsifying the cosmological constant picture. While present observational constraints on the background evolution of the Universe seem to be fully consistent with a static behavior of the DE density, as expected for a standard cosmological constant, complementary and more sensitive tests of a possible time evolution of the DE field might arise from the study of structure formation processes, and in particular from the time evolution of density perturbations at different scales.  In this context, the detection of unexpectedly massive clusters at high redshift \\citep{Mullis_etal_2005, Bremer_etal_2006, Jee_etal_2009,Rosati_etal_2009,Brodwin_etal_2010, Jee_etal_2011,Foley_etal_2011} has recently attracted significant interest as a possible indication of a failure of the standard $\\Lambda $CDM cosmological model. In fact, starting from a Gaussian field of primordial scalar perturbations with an amplitude consistent with the most recent constraints from observations of the Cosmic Microwave Background (CMB) radiation \\citep{wmap7}, the growth of density fluctuations for a $\\Lambda $CDM cosmology predicts an evolution of the halo number density that would imply an extremely low probability for the detection of such clusters. In particular, for the most extreme case represented by the XMMU J2235.3-2557 cluster \\citep{Jee_etal_2009,Rosati_etal_2009} with an estimated mass of $M_{324} = (6.4 \\pm 1.2) \\times 10^{14} M_{\\odot}$ at $z\\sim 1.4$, a $2-3\\sigma $ discrepancy with respect to the fiducial $\\Lambda $CDM model has been claimed by several authors \\citep[see \\eg][]{Jee_etal_2009,Jimenez_Verde_2009,Holz_Perlmutter_2010, Hoyle_Jimenez_Verde_2011, Jee_etal_2011}, although more conservative interpretations have been provided by other analysis \\citep[as \\eg by][]{Mortonson_Hu_Huterer_2011}. The main difficulties in driving sufficiently robust conclusions on a possible tension with the standard $\\Lambda $CDM model from such detections arise as a consequence of the observational uncertainties on the mass determination of the clusters and on the estimation of the statistical significance of the cosmological volumes covered by the cluster surveys \\citep{Sheth_Diaferio_2011,Waizmann_Ettori_Moscardini_2011}. Such significance will need to be better clarified before claiming a discrepancy with respect to the predictions of the standard model. Nevertheless, it is in any case interesting to investigate whether such possible discrepancy could be alleviated by alternative cosmological scenarios.  Several attempts have been made in this direction in the recent past. On one side, a deviation from statistical Gaussianity in the primordial density field has been shown to give rise to a higher number density of massive halos as compared to the standard Gaussian case \\citep[see \\eg][]{Matarrese_Verde_Jimenez_2000, Grossi_etal_2007,Jimenez_Verde_2009,Holz_Perlmutter_2010,Cayon_Gordon_Silk_2010,Sartoris_etal_2010,Hoyle_Jimenez_Verde_2011,LoVerde_Smith_2010}. However, in order to significantly increase the probability of detection this approach requires a level of primordial non-Gaussianity that seems to be already ruled out by CMB constraints \\citep{wmap7}, unless a strongly scale-dependent non-Gaussianity is invoked \\citep[as proposed by \\eg][]{LoVerde_etal_2008,Cayon_Gordon_Silk_2010,Hoyle_Jimenez_Verde_2011}. Alternatively, \\citet{Baldi_Pettorino_2011} have shown that interacting DE models, which are generically characterized by a faster growth of density perturbations with respect to $\\Lambda $CDM due to the presence of a fifth-force mediated by the DE scalar field, also predict a larger number density of massive halos at all redshifts that would result in a higher detection probability.  Both these approaches, however, are faced by the additional problem of overpredicting the abundance of massive clusters at low redshift. In fact, while the detection of massive clusters at high redshifts seems to imply an anomalous growth of density perturbations as compared to the predictions of the standard model, the observed number counts at low redshifts are found to be in very good agreement with the predictions of $\\Lambda $CDM \\citep[see \\eg][]{Reiprich_Boehringer_2002, Vikhlinin_etal_2009b, Mantz_etal_2010, Rozo_etal_2010}. In other words, a non-Gaussian initial density field or a faster growth of density fluctuations can determine an increased number density of massive halos at high redshifts, thereby alleviating the possible tension with the model arising from the unexpected detection of an exceeding number of such objects, but will then also necessarily imply the further growth of such massive halos from these high redshifts to the present time, resulting in a clear discrepancy with present bounds on the cluster mass function at low redshifts.  In the present study we will show how this problem can be naturally solved by an interacting DE model with a suitable self-interaction potential different from the standard exponential or power-law potentials assumed in previous investigations of coupled DE (cDE) cosmologies. In particular, we will show how the existence of a global minimum in the potential, and the consequent ``bounce\" of the DE field on the cosmological constant barrier, is the key feature to address the existence of exceedingly massive clusters at high redshift without necessarily affecting the halo mass function at the present epoch. \\ \\\\  Although our proposed scenario does not benefit in general of the same scaling properties typical of the exponential and power-law potentials, it  naturally provides a mechanism to enhance the expected number of massive halos up to some high redshift and to subsequently reduce it again to the standard $\\Lambda $CDM prediction at the present time, thereby accounting at the same time for the existence of massive clusters at early epochs and for the observed halo abundance at low redshift. We therefore regard the specific model presented in this work as a toy example of how the dynamical nature of a DE scalar field interacting with cold dark matter (CDM) particles can account for the presence of a high redshift peak in the deviation of the expected cluster number counts from the $\\Lambda $CDM prediction. This feature could not arise in standard gravity models or in non-Gaussian cosmological scenarios, and would therefore represent a ``smoking gun\" for the dynamical nature of DE. In the present work we will illustrate in detail this peculiar behavior, also by means of large N-body simulations of structure formation. \\ \\\\  The paper is organized as follows. In Section~\\ref{sec:cDE} we briefly review the main features of standard cDE models for the background evolution of the universe, and we highlight how bouncing cDE models deviate from this standard and well-known behavior. In Section~\\ref{sec:linear} we investigate the evolution of linear density perturbations in the context of both standard and bouncing cDE scenarios, and we show how the latter can provide a possible explanation for a larger amplitude of density perturbations at high redshift without affecting the normalization of the linear matter power spectrum at the present time. This effect is further investigated for the nonlinear regime of structure formation in Section~\\ref{sec:sims}, where we show by means of large N-body simulations how the real halo mass function evolves in time in each of the models considered, and how bouncing cDE models can predict an enhanced number of massive clusters at high redshift without changing the halo abundance at the present epoch. Finally, in Section~\\ref{sec:concl} we draw our conclusions. ",
        "conclusions": "\\label{sec:concl}  In the present paper we have proposed a new model of interacting dark energy based on a SUGRA self-interaction potential and on a constant coupling between dark energy and CDM particles. The most relevant feature of the SUGRA self-interaction potential for the analysis carried out in this work, is the presence of a global minimum that allows the dark energy scalar field to oscillate, thereby changing its direction of motion during the expansion history of the universe. Such feature makes our SUGRA coupled dark energy scenario significantly different from previously proposed models of dark energy interactions, as it allows an inversion of the energy flow between the two interacting components at relatively recent cosmological epochs.  We have therefore compared the cosmological evolution of standard coupled dark energy models based on an exponential self-interaction potential to our proposed SUGRA cDE scenario, both concerning the background dynamics and the evolution of linear and nonlinear perturbations. For the latter task, we have also made use of large N-body simulations results through the publicly available {\\small CoDECS} halo catalogs that include the specific models of dark energy interaction under investigation. \\ \\\\  Already at the background level, the SUGRA cDE scenario shows a series of interesting features. First of all, the inversion of the direction of motion of the dark energy scalar field happens in our specific model at $z_{\\rm inv}\\sim 6.8$ which is a redshift already relevant for astrophysical processes and for the early stages of nonlinear structure  formation. This inversion of motion corresponds to a ``bounce\" of the dark energy equation of state parameter $w_{\\phi }$ on the cosmological constant barrier $w_{\\phi }=-1$, and imprints a specific pattern in the low-redshift evolution of the Hubble function that might provide a way to directly constrain the model. Furthermore, the inversion of the energy flow between dark energy and CDM particle during cosmic evolution implies, for the specific model presented here, that the CDM particles mass has very similar values at the redshift of the CMB and at the present time, thereby avoiding any mismatch of the cosmological parameters between $z_{\\rm CMB}$ and $z=0$.  The most significant result of the present work, however, concerns the evolution of perturbations and the formation of linear and nonlinear structures in the context of the SUGRA cDE scenario, as this is shown to provide appealing features to account for possible anomalous detections of massive clusters at high redshifts. \\ \\\\  At the linear level, we have demonstrated for the first time by numerically solving the linear perturbations equations that a bouncing coupled dark energy model, and in particular the SUGRA cDE scenario presented here, can determine the same amplitude of linear perturbations as a corresponding $\\Lambda$CDM cosmology both at CMB and at the present time, while featuring a significant deviation from $\\Lambda $CDM at intermediate redshifts. In particular, both the amplitude of scalar perturbations at CMB and the value of $\\sigma _{8}$ at $z=0$ are the same between the standard $\\Lambda $CDM model and our proposed SUGRA cDE scenario, while the latter features a significantly larger amplitude of density perturbations at intermediate redshifts, with a peak in correspondence of the dark energy ``bounce\" at $z_{\\rm inv}$. Therefore, the SUGRA cDE model can be simultaneously consistent with CMB and local constraints on the perturbations amplitude, while allowing for significant freedom at intermediate epochs. \\ \\\\  We have then studied the evolution of the halo mass function as predicted by large N-body simulations that include all the characteristic features of the different interacting dark energy models considered in this work. Consistently with what found in previous studies, at high redshifts ($z\\sim 2.5$) all the coupled dark energy models show a larger number of halos with respect to $\\Lambda $CDM over the whole mass range covered by our numerical sample. The enhancement in the halo number density reaches a factor of $\\sim 20$ at $M\\sim 1.0\\times 10^{14}$ M$_{\\odot}/h$ for the most extreme standard coupled dark energy scenario, while the SUGRA cDE model shows an enhancement of a factor $\\sim 4-5$ at the same mass. Furthermore, the enhancement has a clear mass dependence and increases towards larger masses, suggesting that the effect might be significantly larger at masses not covered by the halo sample of our simulations. The subsequent evolution of the halo mass function clearly displays the fundamental difference between the standard coupled dark energy models based on a monotonic self-interaction potential and the bouncing coupled dark energy scenarios, as the SUGRA cDE proposed here. At $z\\sim 1.6$, in fact, all the models still feature a significant excess of massive halos as compared to $\\Lambda $CDM, although the SUGRA cDE model has reduced its enhancement factor much more significantly than the other coupled dark energy models. Finally, at $z=0$ the SUGRA cDE model fully recovers the standard $\\Lambda $CDM mass function over the whole mass range of our sample, while the other coupled dark energy models still show a large excess of massive halos, especially in the range of very massive galaxy clusters, being therefore potentially in tension with available constraint on the cluster number counts.  This result, which represents the main outcome of the present paper, demonstrates for the first time by means of large and fully self-consistent N-body simulations that coupled dark energy models with a suitable choice of the self-interaction potential can simultaneously account for the detection of anomalously massive clusters at high redshifts and for the observed halo mass function at low redshifts. Such peculiar evolution could not arise in other types of models that have been recently invoked as a possible explanation of the unexpected detection of massive clusters at high redshifts, as \\eg non-Gaussian cosmological scenarios, where an enhanced number density of halos at high redshifts necessarily implies a corresponding enhancement at low redshifts. \\ \\\\  To conclude, we have studied a new class of interacting dark energy cosmologies characterized by a ``bounce\" of the dark energy scalar field that is allowed to invert its direction of motion during the cosmic expansion. We have shown by means of linear and nonlinear numerical treatments that this new class of models could be simultaneously consistent with observational constraints at CMB and at the present epoch, while allowing for significant deviations from the standard $\\Lambda $CDM scenario at intermediate redshifts. In particular, we have shown this class of models to possess the (so far) unique feature of simultaneously accounting for unexpected detections of very massive clusters at high redshifts and for the standard cluster abundance at the present time. Such behavior is a direct consequence of a non-trivial dynamics of dark energy at relatively recent cosmological epochs, and therefore represents a specific observational signature of a possible dynamical origin of the accelerated expansion of the Universe."
    },
    "1107/1107.2071.txt": {
        "abstract": "We have used the \\textit{Spitzer Space Telescope} Infrared Spectrograph (IRS) to observe the 5-37~\\micron{} thermal emission of comet 73P/Schwassmann-Wachmann 3 (SW3), components B and C. We obtained low spectral resolution (R$\\sim$100) data over the entire wavelength interval, along with images at 16 and 22 ~\\micron{}. These observations provided an unprecedented opportunity to study nearly pristine material from the surface and what was until recently the interior of an ecliptic comet - cometary surface having experienced only two prior perihelion passages, and including material that was totally fresh. The spectra were modeled using a variety of mineral types including both amorphous and crystalline components. We find that the degree of silicate crystallinity, $\\sim$35\\%, is somewhat lower than most other comets with strong emission features, while its abundance of amorphous carbon is higher. Both suggest that SW3 is among the most chemically primitive solar system objects yet studied in detail, and that it formed earlier or farther from the sun than the bulk of the comets studied so far. The similar dust compositions of the two fragments suggests that these are not mineralogically heterogeneous, but rather uniform throughout their volumes. The best-fit particle size distribution for SW3B has a form dn/da$\\sim$a$^{-3.5}$, close to that expected for dust in collisional equilibrium, while that for SW3C has dn/da$\\sim$a$^{-4.0}$, as seen mostly in active comets with strong directed jets such as C/1995 O1 Hale-Bopp. The total mass of dust in the comae plus nearby tail, extrapolated from to the field of view of the IRS peakup image arrays, is 3-5 x 10$^{8}$ kg for B and 7-9 x 10$^{8}$ kg for C. Atomic abundances derived from the spectral models indicates a depletion of O compared to solar photospheric values, despite the inclusion of water ice and gas in the models. Atomic C may be solar or slightly sub-solar, but its abundance is complicated by the potential contribution of spectrally featureless mineral species to the portion of the spectra most sensitive to the derication of the C abundance.  We find a relatively high bolometric albedo, $\\sim$0.13 for the dust, considering the large amount of dark carbonaceous material, but consistent with the presence of abundant small particles and strong emission features.    %{\\it Facilities:} \\facility{SST (IRS)}. ",
        "introduction": "A major goal of planetary science is to develop a theoretical framework that is capable of reproducing the characteristics of our solar system today. Comets form a direct link to the physical and chemical processing that occurred when our solar system was was only a few Myr old (0.02\\% of its present age). These primitive materials consist of ices of abundant volatile molecules (water, methane, ammonia, carbon monoxide, carbon dioxide), and dust - highly refractory inorganic materials (rock-forming minerals), as well as organic materials of generally intermediate volatility. Of particular interest are the mineralogy, size distribution, and degree of crystallinity of the dust. These provide key information about the condensation of mineral types from the young solar nebula, the  dust-gas chemistry, as well as possible large-scale transport within the nebula.  \\subsection{Origin of Comets and their Dust Material}  Comets fall into two broad classes, based on their orbital dynamics. The nearly isotropic comets (NICs, previously called dynamically-new comets and Oort cloud comets) are those entering the inner solar system from nearly isotropic directions on highly eccentric orbits. They are thought to have been formed and ejected from the Jupiter-Uranus region, and are currently stored in the Oort cloud. The ecliptic comets (ECs), by contrast, are those with low orbital eccentricities  and with inclinations close to the ecliptic, and which include objects significantly affected by the gravity of Jupiter, so that they are in resonant or near resonant orbits with the giant planet, the Jupiter-family comets (JFCs). They are commonly believed to have formed further away from the Sun than the NICs were, near the outer edge of the nebula/disk, and they might be expected to possess distinctly different compositions than the Oort cloud objects. Dynamical studies suggest that the ECs as well as many of the intermediate Halley-family  comets (periodic non-ECs) and possibly some  NICs come from the trans-Neptunian scattered disk, a population still exhibiting these eccentricity and inclination patterns \\citep{dl97,fgb04,lddg06,vm08}. Simulations of the orbital evolution of the planets indicate that significant migration (and possibly even positional swapping by planets \\citep{tsiganis05}) may have occurred during the time of comet formation, so that the distinction between the location of formation of the two comet populations may be blurred.  The composition of the minerals that make up cometary  dust will be a mixture of pre-solar grains that are unchanged remnants of  interstellar medium (ISM), material that accreted into  the molecular cloud that the Sun formed in, materials that condensed within the cloud and proto-solar  nebula itself, and material which has undergone a sequence of condensation and evaporation and annealing  in the circulating proto-planetary disk formed after the nebula collapsed to form the young Sun and its accretion disk. The material that condenses (or re-condenses) in the nebula will be subject to a complex set of gas-grain (and gas-gas) interactions which can produce a radial stratification in the condensation sequence of minerals \\citep{gail98}. Because of this, it is possible that comets forming at different radial distances may have inherited different mineralogies.  It is known that the grains in the interstellar medium have a very low degree of crystallinity, only $\\sim$1\\% \\citep{ciska04,ciska05}. This contrasts to the much higher values observed in comets - exceeding 50\\% in some cases \\citep{har02,har04,lisse07}. Where, when, and how this transformation occurs is still debated, but annealing and sublimation/condensation in the hottest part of the solar nebula is implicated. Spatially resolved spectroscopy of the silicate band of dust in young stellar disk systems indicates that the inner few AU have larger and more crystalline silicates than the more distant regions \\citep{vb03}, suggesting production in the inner regions and radial transport outward.  A related aspect of the formation and processing of dust in the solar nebula is the availability of elemental species. Even the most ``primitive'' meteorites, the CI carbonaceous chondrites, are depleted in carbon and oxygen compared to solar photospheric values \\citep{lodders09a,lodders09b}.  A similar pattern is seen in the mass spectrometer data on 1P/Halley obtained with the \\textit{Giotto} and \\textit{VEGA} spacecraft \\citep{langevin87,jck88,jk91,lawler89}, which presumably sampled material condensed further from the Sun than the meteorites. These are critical elements for the formation of silicates, ices, and organics materials in the solar system. A radial gradient in the depletion pattern might provide clues about the condensation processes in the early solar nebula, and signal a lack of significant radial mixing.  If there is a radial gradient in the characteristics of grain mineralogy and degree of crystallinity of the material in the early solar nebula, it would be manifested in differences between the grain characteristics of these two comet populations. Furthermore, if the grain characteristics of the comet-forming regions evolves while the cometesimals are growing and assembling, these bodies may be internally heterogeneous and possibly stratified. These differences should manifest themselves in differences among the grain characteristics of the comet populations. However, after many passages through the inner solar system, the surfaces of EC comets will become de-volatilized, processed, and may no longer represent the material they formed from. Furthermore, if the mineralogy and degree of crystallinity in the comet-building region was evolving as the comet nuclei were being assembled, then detection of inhomogeneities within the nucleus would provide a window on the temporal changes in the comet-building zone. Any stratification in grain properties would be difficult to detect from a mere snapshot study of its current (and processed) surface.  For these reasons it is important to sample the \\textit{inside} of comet nuclei, particularly those of the ECs. One method is to artificially excavate beneath the surface layers, as was done by the \\textit{Deep Impact (DI)} mission to comet 9P/Tempel 1 (hereafter T1). Another approach is to study comets that have recently fragmented, exposing their more pristine interiors for the first time.  During the past decade or so, a number of ECs have been observed to fragment: D/1993 Shoemaker-Levy 9, 16P/Brooks 2, 51P/Harrington, 57P/du Toit-Neujmin-Delporte, 73P/Schwassmann-Wachmann 3 (hereafter SW3), 101P/Chernykh, 102P/Shoemaker-Holt 1, 174P/Echeclus, 205P/Giacobini, and P/2004 V5 (LINEAR-Hill) \\citep{fernandez09}. Because disrupted comets may be responsible for the production of most of the dust in the zodiacal cloud \\citep{nesvorny10}, understanding the physical and chemical properties of these comets  is important. As of 2006 there were no mid-IR spectra of a fragmented EC suitable for the necessary analysis. However SW3 survived a fragmentation event in 1995, and two orbits later was poised for a close approach to the Earth. This provided an unprecedented opportunity to study one of these objects in detail.  In this paper, we describe observations of the B and C fragments of SW3 using the Infrared Spectrograph (IRS) of the \\textit{Spitzer Space Telescope}. Although \\textit{Spitzer} was more distant from the comet than the Earth at its closest approach, its superior sensitivity, wavelength coverage, and lack of telluric spectral contamination provided us with the best opportunity to examine this important object in the mid-IR. Additional supporting ground-based visual, near-IR, and mid-IR observations that were useful for interpreting the \\textit{Spitzer} data are described in the Appendix. We will also compare the results of the analysis of the \\textit{Spitzer} observations with other independent mid-IR observations of SW3 \\citep{har10} as well as other comets studied spectroscopically and using \\textit{in situ} sampling techniques.  In this study we will be able to:  $\\bullet$ obtain information on the mineral content and degree of crystallinity on (potentially) the most pristine cometary material yet measured through remote sensing  $\\bullet$ compare the abundances of the various minerals between two cometary fragments that represent different fractions of material formerly buried below the surface of the nucleus  $\\bullet$ compare the atomic abundances of SW3 with other comets studies both spectroscopically and \\textit{in sutu}, primitive meteorites, and the Sun, to investigate the pattern of elemental depletions among these solar system objects ",
        "conclusions": "The inner solar nebula possessed the thermal and chemical ingredients for the production of silicate material significantly more crystalline than the raw interstellar dust from which it formed. Gas-gas and gas-grain chemical reactions likely converted some portion of the primordial amorphous carbon in ISM grains into a variety of carbon-containing molecules. Both the crystalline material and the amorphous C-depleted and hydrocarbon molecule-rich material were mixed radially, although cometary material formed very early or far from the inner nebula would not be affected to the degree of those formed later or closer to the sun. Spectroscopic observations of both B and C components of 73P/Schwassmann-Wachmann 3 reveal a low degree of grain crystallinity compare to other comets (but higher than the ISM), higher amorphous carbon content than most other comets, and depletion of most C-rich molecules relative to other comets. These are all indicative of being more primitive than other comets studied remotely, and suggest it was formed earlier than most, or further from the warm inner nebula, or both.  No significant difference in the crystallinity nor amorphous carbon content was observed between the two components in either the \\textit{Spitzer} data nor the independently-observed and modeled Michelle spectra by \\citet{har10}. The difference in the rates of nucleus fragmentation suggests that we are likely digging deeper into the pre-fragmentation nucleus in B than in C. Despite this, they appear spectroscopically very similar, arguing against significant heterogeneity and/or layering.  One of the most significant differences observed, the relative strength of the water vapor bands, may not point to any intrinsic difference in the two components, but instead to their being significant temporal changes in their water vapor production, or significant differences in the distribution of water ice sources in the extraction beam.  All comets studied spectroscopically exhibit a general under-abundance of O-rich species. It is likely that much of the O is tied up in grains that are spectroscopically featureless, or at least not easily detectable with current techniques. These do not, however, explain the deficit in the O abundance when measured \\textit{in situ} (such as Halley), which is likely dominated by the exclusion of the O on the form of H$_{2}$O gas. Some of it may be in the disk gas that that remains at high altitude in the proto-solar nebula as the dust from which comets form settle to the mid-plane, as has been seen in other  protostellar disks.  A better understanding of the chemical and physical evolution of the early solar system through the analysis of cometary material will require many studies using a variety of complementary techniques, both remote sensing (spectra, imaging) and \\textit{in situ} sampling (dust fluence experiments, sample returns, remote sample analysis)."
    },
    "1107/1107.0807_arXiv.txt": {
        "abstract": " ",
        "introduction": "Pulsars are rotating neutron stars and have a strong magnetic field, by which they emit optical, radio , X-ray and $\\gamma$-ray pulses.   Although it is not well understood yet about where or how the pulsar emission takes place, the magnetic dipole model of pulsars simply states, from the viewpoint of energetics,  how the rotation energy of neutron stars is converted into pulsar emission \\cite{SLS}. First of all, let us briefly see the physical mechanism. Consider the sphere with radius $a$ composed of uniform magnetization ${\\bf M}$ (${\\bf M}//{\\hat z}$). Then the solution of the vector potential is \\beq A_{\\phi}=\\frac{4\\pi}{3}Ma^2\\frac{r_<}{r_>}\\sin\\theta, \\eeq where $(r_>, r_<)$ are the larger and smaller of $(r, R)$ \\cite{jac}. The magnetic field ${\\bf B}$ is then given by \\beqa B_r&=&\\frac{1}{r\\sin\\theta}\\frac{\\partial}{\\partial\\theta}(\\sin\\theta A_\\phi)\\nonumber\\\\ B_\\theta&=&-\\frac{1}{r}\\frac{\\partial}{\\partial r}(rA_\\phi) \\nonumber\\\\ B_{\\phi}&=& 0. \\eeqa The magnetic field outside the sphere is \\beqa B^{({\\rm out})}_\\theta&=&\\frac{|{\\bf m}|\\sin\\theta}{r^3}\\nonumber\\\\ B^{({\\rm out})}_r&=&\\frac{2|{\\bf m}|\\cos\\theta}{r^3}. \\label{bout} \\eeqa Hence it takes the maximum magnitude at $\\theta=0$, \\beq B^{({\\rm out})}_{\\rm max}=2|{\\bf m}|r^{-3}, \\label{maxb} \\eeq with the total magnetic moment ${\\bf m}=4\\pi/3a^3 {\\bf M}$. When we consider a star with radius $R$ ($a=R$), $|{\\bf m}|$ can be related to the magnetic field at the magnetic pole of the star, $B_p\\equiv 2|{\\bf m}|R^{-3}$, by way of Eq.~(\\ref{maxb}). Accordingly the magnetic-field lines can be drawn by solving the differential equation, \\beq \\frac{dr}{B^{({\\rm out})}_r}=\\frac{rd\\theta}{B^{({\\rm out})}_\\theta} \\eeq to give \\beq r=R(\\sin\\theta/\\sin\\theta_0)^2, \\eeq outside the star (Fig.~\\ref{dipole}).  \\begin{figure}[h] \\begin{minipage}{0.4\\textwidth} \\begin{center} \\includegraphics[width=4.8cm]{mag.eps} \\end{center} \\caption{Dipole magnetic field-lines outside the star.} \\label{dipole} \\end{minipage} \\hspace{\\fill} \\begin{minipage}{0.56\\textwidth} \\begin{center} \\includegraphics[width=4.5cm]{emission.eps} \\end{center} \\caption{Angular distribution of the luminosity.} \\label{emission} \\end{minipage} \\end{figure}  The magnetic field is uniform inside the star, \\beq \\bB^{({\\rm in})}=\\frac{2\\bf m}{R^3}. \\label{bin} \\eeq  The {\\it light cylinder} is defined as a cylinder with an axis along the pulsar rotation axis and with the radius $r_c=c/\\omega$, where the velocity of the co-rotating frame with the pulsar attains the speed of light. Then the emission activity of the pulsar has been considered to originate from the magnetosphere with the radius $r_c$, while its mechanism is not clear yet \\cite{SLS}.  Denoting the inclined angle of the magnetic moment $\\bf m$ to the rotation axis ${\\bf e}_3$ by $\\alpha$, we can evaluate the luminosity (energy radiation per unit time) $L_m$. Then we can immediately see that the component perpendicular to ${\\bf e}_3$, ${\\bf m}_\\perp$, is responsible to the luminosity, which is rotating in the plane perpendicular to the rotation axis, ${\\bf m}_\\perp=|{\\bf m}\\sin\\alpha|[{\\bf e}_1\\cos(\\omega t)+{\\bf e}_2\\sin(\\omega t)]$, with ${\\bf e}_i\\cdot{\\bf e}_j=\\delta_{ij}$. Thus the luminosity becomes equivalent with the one brought about by two oscillating magnetic dipoles (see Fig.~\\ref{emission}). A standard formula gives \\beq \\frac{d L_m}{d\\Omega_\\omega}=\\frac{\\omega^4}{4\\pi c^3}|{\\bf m}\\sin\\alpha|^2\\frac{1}{2}(1+\\cos^2\\theta_\\omega) \\eeq for differential luminosity, where $\\theta_\\omega$ is an angle relative to the rotation axis \\cite{jac}.  Finally we have \\beqa L_m&=&\\int d\\Omega_\\omega \\frac{d L_m}{d\\Omega_\\omega}\\nonumber\\\\ &=&\\frac{2\\omega^4}{3c^3}|{\\bf m}\\sin\\alpha|^2\\nonumber\\\\ &=&\\frac{B_p^2R^6\\omega^4\\sin^2\\alpha}{6c^3}, \\label{lumi} \\eeqa for total luminosity. It would be interesting here to see that the luminosity (\\ref{lumi}) can be recast into the form, \\beq L_m\\simeq c\\left(\\frac{B_pR^3}{2r_c^3}\\right)^2 r_c^2\\sim c\\frac{B(r_c)^2}{8\\pi}\\cdot 4\\pi r_c^2, \\eeq which implies that the energy density of the magnetic field ($\\sim B(r_c)^2/8\\pi$) flows outward through the surface of $4\\pi r_c^2$ with the velocity of light $c$.  Assuming that this energy release originates from the rotational energy of neutron star, \\beq \\frac{d E_{\\rm rot}}{dt}=-L_m, \\label{rote} \\eeq where the rotation energy is given as $ E_{\\rm rot}=\\frac{1}{2}I\\omega^2 $ with the moment of inertia $I$, assuming the rigid body. From Eqs,~(\\ref{lumi}), (\\ref{rote}), we have \\beq P{\\dot P}=\\frac{2\\pi^2 R^6\\sin^2\\alpha}{3Ic^3}B_p^2, \\eeq where $P=2\\pi/\\omega$ is the period of the rotation. The {\\it characteristic age} $\\tau_c$ is defined by \\beq \\tau_c\\equiv P/2{\\dot P}, \\label{age} \\eeq which gives a rough estimate of the age of pulsars.   The spin-down rate $\\dot \\omega$ is given as ${\\dot\\omega}\\propto -\\omega^3$ from Eqs.~(\\ref{lumi}) and (\\ref{rote}), so that the {\\it braking index} $n\\equiv -\\omega{\\ddot\\omega}/({\\dot\\omega})^2$ is three in this model. Note that the observed values of the braking index are less than three, e.g. $n=2.515\\pm 0.005$ for the Crab, and it is usually considered that the significant fraction of the rotation energy is carried away by a pulsar wind \\cite{bec}. Thus the magnetic dipole model is not sufficient to understand the emission processes in pulsars, while it can provide plausible estimates for the magnetic field and the characteristic ages.  Nowadays observation of $P-{\\dot P}$ for pulsars are summarized on the plane \\cite{bec}, where we can see three clusters of pulsars: ordinary radio pulsars, the millisecond pulsars and {\\it magnetars}. We can easily see that most of the radio pulsars are centered around $10^{12}$G.  Magnetars are compact stars with huge magnetic field ($B=O(10^{14-15})$G), including soft $\\gamma$-ray repeaters (SGRs) or anomalous X-ray pulsars (AXPs)\\cite{woo,har,mere}. It exceeds the critical magnetic field defined as \\beq B_{\\rm QED}=\\frac{m_e^2c^3}{e\\hbar}\\simeq 4.4\\times 10^{13}{\\rm G}, \\eeq which is obtained by equating the spacing of the Landau levels, $\\delta E_e=\\hbar eB/(m_e c)$, with the mass, $m_ec^2$. So it presents a criterion about how the relativistic effects are important and the quantum effects such as $e^+e^-$ creation or photon splitting becomes significant. It should be also interesting to recall the relation, $1{\\rm MeV}^2\\simeq 10^{13}$G in the natural unit. So, $B=10^{15}$G expected in magnetars may correspond to $10$MeV, which gets to the strong-interaction energy scale or the Fermi energy of nuclear matter or quark matter at nuclear density.  We can easily asses how the magnetic field affects EOS inside compact stars. As shown in Eq.~(\\ref{bin}) we can consider the uniform magnetic field inside the stars. Then $\\delta E_p$ is small for protons, $\\delta E_p\\ll m_p$, while it is comparable with light quark mass, $\\delta E_q\\sim m_q$, for $B=O(10^{15})$G. On the other hand, the magnetic-interaction energy can be estimated by a simple formula, $E_{\\rm mag}=\\mu_i B$, with the Dirac magnetic moment, $\\mu_i=e_i/(2m_i)$, for $i$-th particle having mass $m_i$ and electric charge $e_i$. For electrons it gives $O(\\rm keV)$ for the canonical value of $B=O(10^{12})$G, which is comparable with the atomic-energy scale. So we may  easily expect that thermodynamic quantities or EOS of the neutron-star envelope or the crust region should be much affected by the magnetic field \\cite{hae}. On the other hand, it becomes very tiny for protons ($E_{\\rm mag}\\sim {\\rm keV}$) even for $B=O(10^{15})$G, which implies that the magnetic field little affect the EOS of nuclear matter. Since it amounts to $E_{\\rm mag}\\sim{\\rm MeV}$ for light quarks, the magnetic field looks to modify the EOS of quark matter. However, the Fermi energy should be very large ($O(100) {\\rm MeV} \\gg \\delta E_q$), and a large number of the Landau levels are occupied, which may be well approximated by the usual Fermi sphere. Thus we can conclude the effect of the magnetic field little modifies the EOS in the core region, even for $B\\sim 10^{15}$G.  Other high-field radio emitting pulsars with $B\\ge 10^{14}$G or rotation-powered pulsar/magnetar transition objects have been also observed \\cite{gav}. These observations may give a hint about the relationship between magnetic activity and neutron star spin-down.   The origin of such strong magnetic field has been a basic but a long-standing problem since the first discovery of pulsars, while many people believe that it originates from dynamo scenario due to the charged current or inheritance of the magnetic field from the progenitor main-sequence stars (fossil-field hypothesis) \\cite{chan}. This problem becomes a current issue, stimulated by the discoveries of magnetars. Due to the dynamo scenario, magnetic field is produced or maintained by the rotation or convection of charged fluid, but seems irrelevant for generation of high magnetic field in magnetars. It requires that magnetars be born with very short rotation periods of the order of 1-2 ms, and may not be supported by observations by two reasons \\cite{mere}. First, The combination of high magnetic field and very rapid rotation is expected to impart a high velocity to the neutron star. However, up to now, the observational evidence for large spatial velocities in SGRs and AXPs is poor. Secondly, a large fraction of the rotational energy of a newly born magnetar, a few $10^{52}$erg, is lost due to the strong magnetic braking. However, an estimate of the explosion energy of the remnants containing magnetars yields values close to the canonical supernova explosion energy of $10^{51}$erg, implying initial periods longer than 5 ms.  Fossil-field hypothesis assumes the conservation of magnetic flux during the evolution of main-sequence progenitor star to a compact star: $B=(R_M/R)^2B_M$ with $R_M,B_M$ being the radius and magnetic field of the progenitor, respectively. It should be interesting to compare the radius $R$ with . Taking the sun as a typical main-sequence star, $(B_\\odot)_{{\\rm Max}}\\simeq {\\rm several~thousands}$~G and $R_\\odot\\simeq 7\\times 10^{10}$cm. Squeezing the magnetic flux from $R_\\odot$ to $R$, we have $R=R_\\odot(B_\\odot/B)^{1/2}\\simeq 10$km for usual pulsars with $B=10^{12}$G, which should be consistent with standard neutron stars. However, $R<{\\rm several}\\times 10^{4-5}$cm for $B\\simeq 10^{14-15}$G, which should be compared with the Schwartzschild radius, $R_S=2GM/c^2=3(M/M_\\odot)\\simeq 3\\times 10^5$cm for $M \\simeq M_\\odot$. Thus we can see that $R$ should be comparable with $R_s$ and may fall below it in the extreme case, if the fossil-field hypothesis is applied to the generation of the magnetic field in magnetars.  There is another possibility: a microscopic origin due to ferromagnetism or spin polarization of hadron matter \\cite{nov}. Actually Makishima suggested the hadronic origin of the magnetic field in binary X-ray pulsars or radio pulsars \\cite{mak}, since the accumulation of the observational data shows a peak with a narrow width. Sometimes high-density nuclear matter is expected to show a resemblance with $^3$He system, by the scaling argument\\cite{tam}, where quantum effects are essential and it shows the $^3P$-type superfluidity in some physical conditions.  Correspondingly neutron matter has been shown to exhibit the $^3P_2$ superfluidity around the nuclear density \\cite{tam}. It is also well-known that $^3$He is very close to a ferromagnetic state under the high pressure due to the large magnetic susceptibility. Microscopic calculations of nuclear matter have been repeatedly performed to find out a possibility of spontaneous spin polarization inside pulsars, but the negative results have been reported so far \\cite{mat}. The magnetic susceptibility of nuclear matter increases monotonously with density, so that the ferromagnetic phase is not be expected in nuclear matter at any density. At high-density we can also expect quark matter and we may be tempted to ask the possibility in quark matter. Nowadays there have been actively studied about the possible evidences of hadron-quark transition in astrophysical phenomena as well as relativistic heavy-ion collisions \\cite{bec,R-star2,sag}. If such phase transition occurs inside compact stars, it may affect thermal and magnetic properties as well as equation of state (EOS). Here we ask a microscopic origin of the magnetic field by considering uniform quark matter inside compact stars \\cite{tat00}.  We consider the possibility of spontaneous spin polarization by the use of QCD. First, the evaluation of the total energy for spin-polarized quark matter has been done by using the one-gluon-exchange (OGE) interaction at zero temperature \\cite{tat00,nie05,pal09}. Subsequently, the magnetic susceptibility has been studied within the Landau Fermi-liquid theory \\cite{bay}, taking into account the screening effect for gluon propagation \\cite{tat082,tat083,pal10}. Finally we present a magnetic phase diagram on the density-temperature plane.  We also intend to address another interesting magnetic aspect of quark matter, which is characterized by the spatial modulation of the magnetic moment \\cite{DCDW}. It is an inhomogeneous phase accompanying the chiral transition. In the standard scenario, the quark condensate, which is the order parameter of spontaneous symmetry breaking (SSB) of chiral symmetry, decreases due to the Pauli principle as density is increased. It is eventually vanished at some density-temperature point, which means the restoration of chiral symmetry. However, this may not be a unique scenario. Recall the FFLO state in the superconductivity; it has a spatially non-uniform order parameter and is considered to appear in the vicinity of the critical point, when two Fermi spheres with different spins have different sizes \\cite{fflo}. Recent studies have shown that there possibly appear various kinds of inhomogeneous phases in the vicinity of the chiral-symmetry restoration \\cite{bas,nic,brin}. Among them we consider a special one called dual-chiral-density-wave (DCDW) state in detail, because it exhibits many interesting theoretical features and is expected to bear various implications in the light of compact-star phenomena.  Nowadays there have been done many theoretical studies about the deconfinement transition in high-density nuclear matter by using the MIT bag model \\cite{mit} or other effective models of QCD to find the EOS including quark degrees of freedom. Unfortunately the lattice simulation is not possible, because the numerical calculation  suffers from the so called sign problem \\cite{R-lat}. Thus there are ambiguities about the properties and the critical density of the deconfinement transition. Here we only assume the presence of quark matter without resource to EOS and the details of the deconfinement transition. ",
        "conclusions": "It should be important to study EOS at high-density region from various viewpoints. Here we tried to extract information about the properties of hadron or quark matter and EOS, confronting the magnetic phenomena of pulsars. We have seen two kinds of the magnetic properties in quark matter: one is spontaneous spin polarization and the other is the spin density wave. The former is similar to the itinerant electrons in QED, while the latter is related to the chiral symmetry of QCD. These discussions may be in a rather primitive level and must be verified by observations or more elaborate theoretical studies toward more realistic description. Comprehensive study of these magnetic properties is also needed to study their competition.  Energetics or the mechanism of the giant flares observed in some magnetars may give a direct hint about the origin of the magnetic field. Thermal evolution of compact stars should be important for us to get the information about the EOS and properties of high-density matter. Actually the magnetic properties is not directly related to EOS but may manifest through the thermal activities of compact stars.  There have been done many works about color superconductivity (CSC) in high-density matter \\cite{csc}. It also has implications on compact star phenomena. So it should be interesting to study the interplay of CSC and magnetic properties in quark matter. In ref.\\cite{nak03} we have discussed a coexistence of CSC and ferromagnetic order.  More elaborate studies are necessary for inhomogeneous phases: relations among the various phase should be figured out as well as their properties. For example, there seem many resemblances between pion condensation in hadron matter \\cite{pio} and DCDW state. One may find a kind of hadron-quark continuity across the deconfinement transition. They may also appear during relativistic heavy-ion collisions. So it is interesting to consider how we can observe them in this context."
    },
    "1107/1107.0200_arXiv.txt": {
        "abstract": "Significant advances have been made in the understanding of the  diffuse Galactic hard X-ray  continuum emission using data from the \\INTEGRAL\\ observatory. The diffuse hard power-law component seen with the \\INTEGRALSPI\\ spectrometer has been identified with inverse-Compton emission from relativistic (GeV) electrons on the cosmic microwave background and Galactic interstellar radiation field. In the present analysis, SPI data from 2003 to 2009, with a total exposure time of $\\sim 10^{8}$ s, are used to derive the Galactic ridge hard X-ray spatial distribution and spectrum between 20 keV and 2.4 MeV. Both are consistent with predictions from the \\GALPROP\\ code. The good agreement between measured and predicted emission from keV to GeV energies suggests that the correct production mechanisms have been identified. We discuss the potential of the SPI data to provide an indirect probe of the interstellar cosmic-ray electron distribution, in particular for energies below  a few GeV. ",
        "introduction": "The Galactic ridge is known to be an intense source of continuum hard X-ray and \\gray{} emission. The hard X-ray emission was discovered by a rocket experiment in 1972 \\citep{Bleach72}, and interstellar emission has subsequently been observed from keV to MeV energies by the satellites HEAO-1, EXOSAT, Temna, ASCA, Ginga, RXTE,  CGRO/COMPTEL, GRANAT/SIGMA, and more recently by Chandra, XMM/Newton, along with INTEGRAL. Previous analyses of INTEGRAL data showed that, up to 100 keV, a large fraction of the total emission from the inner Galaxy is due to point sources, and the diffuse emission associated with the Galactic ridge is only one tenth of the total Galactic emission in the 25--100 keV band \\citep{Lebrun04, Strong04, Bouchet05, Bouchet08}. These analyses also showed that the diffuse emission dominates the hard X-ray sky above 300 keV.  Continuum emission of a diffuse, interstellar nature is expected in the hard X-ray and \\gray{} regime from several physical processes: positron annihilation (through intermediate formation of positronium), inverse-Compton (IC) scattering of the interstellar radiation field (ISRF) and bremsstrahlung on the interstellar gas from cosmic-ray (CR) electrons and positrons, and via decay of neutral pions produced in interactions of CR nuclei with the interstellar gas. Extensive studies of the Galactic diffuse \\gray{} emission in the context of CR propagation models have been carried out by, e.g., \\citet{Strong00, Strong04, Strong10} and \\citet{Strong2010}.  In the present analysis, the data accumulated by the SPI spectrometer onboard the \\INTEGRAL\\ observatory are used to derive the spatial morphology and the spectral shape of the Galactic diffuse emission, taking advantage of the greatly increased observational data and significant advances in the analysis techniques. This builds on our previous work, described in \\citet{Bouchet08} and \\citet{Porter08}, where we have presented sky maps and spectra of the Galactic plane and demonstrated the presence of a hard power-law continuum emission, which was interpreted as IC emission from CR electrons and positrons upscattering the Galactic ISRF. For further introductory material and background information on this topic we refer the reader to these earlier papers. ",
        "conclusions": "New results on the diffuse emission spatial morphology and spectrum in the range 20~keV--2.4~MeV have been obtained from the analysis of 6 years of \\INTEGRALSPI\\ data. Over this energy range, what is seen as \\diffuse emission is the result of the superposition of several physical processes; annihilation radiation, cosmic nuclear \\gray{} lines, diffuse continuum due to interstellar emission, and unresolved sources. In the present analysis we have been able to isolate each of these diffuse components. We have explored uncertainties and limitations due to the data reduction along with the spatial morphology modelling.  The \\diffuse emission intensity in the central radian ($\\vert l \\vert < 30^{\\circ}$, $\\vert b \\vert < 15^{\\circ}$) is estimated to be one tenth of the total emission (including sources) below 100~keV and one third in 100--300~keV band. The \\diffuse emission spectrum has the following main features:  \\begin{itemize}  \\item The \\diffuse continuum spectrum (apart from positronium) is  fitted by a power law of index $1.4 - 1.5$, with a flux at 100~keV of $1.1 \\times 10^{-4}$ \\phn. This power law, thought to be related to interstellar emission, can be decomposed into two spatial components: the IC component with a power law of index 1.8 and a flux at 100~keV of $\\sim 10^{-4}$ \\phn and another component which can be modelled with the an extinction-corrected 4.9 $\\mu$m DIRBE-based template, whose spectral shape is a power law with an index of $\\sim 1$ and a flux at 50~keV of $\\sim 3-4 \\times 10^{-5}$ \\phn. This additional component is weak below 200 keV compared to the IC component.  \\item  The diffuse continuum flux around 1 MeV is found compatible with the \\textit{COMPTEL/CGRO} measurement.  \\item The IC emission distribution predicted by the \\GALPROP\\ code is in fair agreement with the data. However, a model with an electron spectrum increased by a factor 2 over the standard model based on the electrons (plus positrons) measured by \\FermiLAT\\ is in better agreement. Also, an increased ISRF in the Galactic bulge or a large Galactic CR halo are other reasonable possibilities to that can lead to an increased flux. The data analysed in this paper do not allow a distinction to be made between these possibilities.  \\item An additional component is required below 50~keV. This excess over the IC emission is well modelled with the NIR/DIRBE 4.9 $\\mu$m map. This low-energy component has an exponential spectrum with a cutoff at 8~keV and a flux at 50~keV of $\\sim 2\\times 10^{-4}$ \\phn. It can interpreted in terms of the stellar origin as proposed by \\citet{Krivonos07}.  \\end{itemize}  The \\diffuse continuum emission spectrum obtained with the present analysis confirms and improves on the results reported in \\citet{Bouchet08} and \\citet{Porter08}. Below 100~keV, it is also compatible with results obtained from \\INTEGRALIBIS\\ analysis \\citep{Krivonos07, Turler10}. In addition, we found that our global \\diffuse spectrum is consistent through the measured characteristics of nuclear lines ($^{26}$Al, $^{60}$Fe) and annihilation radiation spectrum.  There is no detection in the SPI energy range of the ``Fermi bubbles'' at the present level of sensitivity.  The \\INTEGRALSPI\\ will continue to provide new data for at least the next 3 years. Meanwhile, the \\FermiLAT\\ will be operating at least over this time span, and analysis of its data continue to yield further information on the diffuse \\gray{} emission $\\gtrsim 100$~MeV. An analysis showing the complementarity of the data provided by these instruments enabling us to probe the physical processes producing the diffuse emission is given in \\cite{Strong2010}. Together with existing data from other instruments, the data from these two currently operating missions allows the investigation of the CR electron spectrum at all energies, which will eventually enable an unambiguous decomposition of the diffuse \\gray{} sky."
    },
    "1107/1107.1426_arXiv.txt": {
        "abstract": "{} {Observational data on companion statistics around young stellar systems is needed to flesh out the formation pathways for extrasolar planets and brown dwarfs. Aperture masking is a new technique that is able to address an important part of this discovery space.} {We observed the two debris disk systems HD\\,92945 and HD\\,141569 with sparse aperture masking (SAM), a new mode offered on the NaCo instrument at the VLT. A search for faint companions was performed using a detection strategy based on the analysis of closure phases recovered from interferograms recorded on the Conica camera.} {Our results demonstrate that SAM is a very competitive mode in the field of companion detection.  We obtained $5\\sigma$ high-contrast detection limits at $\\lambda/D$ of $2.5\\times10^{-3}$ ($\\Delta L' = 6.5$) for HD\\,92945 and $4.6\\times10^{-3}$ ($\\Delta L' = 5.8$) for HD\\,141569. According to brown dwarf evolutionary models, our data impose an upper mass boundary for any companion for the two stars to, respectively, 18 and 22 Jupiter masses at minimum separations of 1.5 and 7~\\au. The detection limits is mostly independent of angular separation, until reaching the diffraction limit of the telescope.} {We have placed upper limits on the existence of companions to our target systems that fall close to the planetary mass regime. This demonstrates the potential for SAM mode to contribute to studies of faint companions. We furthermore show that the final dynamic range obtained is directly proportional to the error on the closure phase measurement.  At the present performance levels of 0.28~degree closure phase error, SAM is among the most competitive techniques for recovering companions at scales of one to several times the diffraction limit of the telescope. Further improvements to the detection threshold can be expected with more accurate phase calibration.} ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[scale=.45]{Image_Plane.eps}\\hspace{1cm} \\includegraphics[scale=.45]{Image_FFT.eps} \\caption{Left panel: Diffraction pattern observed on star HD\\,135549. This image is a single frame of 0.4\\,ms integration time, using the 7-hole pupil mask. Background noise is $\\approx$12 ADU, and pic brightness 800 ADU. Right panel: Fourier transform of the diffraction pattern. Each peak corresponds to a baseline pair.} \\label{Image}% \\end{figure*}   Since the first imaging of a debris disk around $\\beta$~Pictoris \\citep{1984Sci...226.1421S}, and the remarkable successes of the radial velocity technique for exoplanet discovery heralded nearly a decade later by a companion to solar-like 51~Pegasi \\citep{1995Natur.378..355M}, the understanding of planetary systems formation and evolution has become one of astronomy's greatest challenges.  Today, at least 22 debris disks have been imaged at optical, infrared, or submillimeter wavelengths.  Such disks are thought to be the cradle where planetary systems have recently formed. The observed dust component is not primordial but instead produced by collisions among larger rocky bodies. This makes debris disks obvious places to search for planets, and many show possible indirect signs of their presence, such as dust structures or disk asymmetries.  The recent detection and confirmation of a $\\sim9$~\\Mjup\\ planet within the disk of \\betapic\\ (semi-major axis between 8 and 15~\\au) using direct deep imaging validates the link between structures in debris disks and the presence of planets \\citep{2009A&A...493L..21L,2010Sci...329...57L}. Besides this planet, only a couple of planetary-mass objects have been detected around stars with a debris disk. In fact, most debris disks have been imaged around young bright A- or F-type stars, which are either rapid rotators and/or active stars. The search for companions with radial velocity techniques around these stars is therefore extremely difficult. Key breakthroughs have recently been made with high-contrast imaging: a $<3$-\\Mjup\\ planetary companion was detected in the outskirts of Fomalhaut's debris disk \\citep[119~\\au\\ from the star;][]{2008Sci...322.1345K}, while four planetary companions of 7, 10, 10, and 10~\\Mjup\\ were imaged at 68, 38, 24, and 14~\\au , respectively, from HR~8799 \\citep{2008Sci...322.1348M,2010Natur.468.1080M}. It is not clear whether these objects could have formed through core accretion -- and not from gravitational instabilities -- the foremost theory of planet formation, although \\betapic~b is a serious candidate for this formation scenario \\citep{2010Sci...329...57L}.  Most companions detected with high-contrast imaging, such as AB~Pic~b \\citep{2005A&A...438L..29C} or 2M1207~b \\citep{2004A&A...425L..29C}, are believed to be either too far or too massive to have formed through core accretion. An alternate pathway gaining widespread interest is a stellar formation mechanism, i.e., formation during the fragmentation and collapse of a molecular cloud. There are many unanswered questions raised by these models. Do these two processes for forming planets and brown dwarfs/low-mass stars operate exclusively (does one inhibit the other)? Can core accretion from a disk operate when a brown-dwarf companion exists at large separation? In contrast, can disks and planets exist around tight binary systems, composed of the primary star and a very close low-mass star or brown dwarf? How do massive companions impact the dynamic evolution of planets and disks?  To answer these questions, it is necessary to probe the innermost regions of debris disks when searching for substellar companions. For the closest debris disk, direct imaging proved to be a successful technique, as shown for $\\beta$~Pic or Fomalhaut. But most of the young associations (as well as stellar forming regions) lie at more than a hundred parsecs. At such a distance, 10 \\au\\ is equivalent to 100\\,mas, less than twice the diffraction limit of an 8-meter telescope in the near-infrared. This observational domain is not filled by differential imaging techniques, whose inner working angle is typically of few resolution elements.  It is neither filled by long baseline interferometry: the field of view of a single mode interferometer is too small, equal to the resolution element of one of its telescopes. The effectiveness of both techniques is reviewed in detail in \\citet{2010A&ARv..18..317A}. In terms of angular separation, aperture masking is the missing link in the field of high dynamic observations. It gives a relatively high dynamic range (5 to 7 mag) at a specific angular separation ($\\approx 100\\,$mas) which is of greatest interest for planetary formation.   In this paper, we present three stars observed with an aperture mask at the VLT.  In the next section, the instrument is described, along with the observational procedure and data reduction methods. In section~\\ref{test}, we present a clearly resolved detection of the binary star HD\\,135549 for illustrative and pedagogical purposes. Section~\\ref{debris} reports on the non detection contrast limits obtained with SAM for the two debris disks HD\\,92945 and HD\\,141569. We present these limits as a function of flux ratio and infer the mass limit derived from current models. ",
        "conclusions": "We have shown that aperture masking gives detection limits of the order of $\\Delta~{\\rm L'\\,mag}=6$, with an inner working angle close to $\\lambda$/2D. These results confirm previous detection limits obtained by the same aperture masking technique on the Keck telescope \\citep{2011ApJ...731....8K,2011ApJ...730L..21H} In terms of scientific impact, this observational domain is important because it corresponds to a few astronomical units at a hundred parsec, the distance where the closest formation regions are. The scientific importance of this parameter space is highlighted by T Cha b detected by the same technique \\citep{2011A&A...528L...7H}.  Simulations of the older debris disk tend to show that a magnitude or two in dynamic range is still needed to observe disk shaping planets. Considering HD\\,141569, for example, \\citet{2009A&A...493..661R} show that the disk geometry could be best modeled by a flyby star and a planet of a few Jupiter masses. For these kinds of objects, a detection limit of two Jupiter masses would require a precision on the closure phases of 0.01 degree, something only  possible if we understand how to precisely account for the systematics errors.  This paper has to be considere along with other direct detection techniques.  In the case of HD\\,92945, \\citet{2007ApJS..173..143B} used a spectral differential technique to derive an upper limit of 10.8 mag at 0.5'' in the H band (5\\,$\\sigma$). This limit assumes a clear methane signature (a spectral type T8 for the planet) which may not be the case. This assumption is not required for a differential technique using sky rotation: in their Gemini survey, \\citet{2007ApJ...670.1367L} obtain the same dynamic range (10.8 mag; 5\\,$\\sigma$) at the slightly larger separation of 0.75''. Overall, differential imaging techniques often give a dynamic range that is higher than SAM, but with larger inner working angles (typically a few $\\lambda$/D).  The capacities of future extreme adaptive optics to observe at smaller inner working angle, compared to the ability for aperture masking to probe at higher dynamic range by understanding the systematics, will decide the fate of SAM with respect to faint companion detection."
    },
    "1107/1107.3086_arXiv.txt": {
        "abstract": "Timing pulses of pulsars has proved to be a most powerful technique useful to a host of research areas in astronomy and physics. Importantly, the precision of this timing is not only affected by radiometer noise, but also by intrinsic pulse shape changes, interstellar medium (ISM) evolution, instrumental distortions, etc. In this paper we review the known causes of pulse shape variations and assess their effect on the precision and accuracy of a single measurement of pulse arrival time with current instrumentation. Throughout this analysis we use the brightest and most precisely timed millisecond pulsar (MSP), PSR~J0437$-$4715, as a case study, and develop a set of diagnostic tools to evaluate profile stability in timing observations. We conclude that most causes of distortion can be either corrected by state-of-the-art techniques or taken into account in the estimation of time-of-arrival (TOA) uncertainties. The advent of a new generation of radio telescopes (e.g. the Square Kilometre Array, SKA), and their increase in collecting area has sparked speculation about the timing precision achievable through increases in gain. Based on our analysis of current data, we predict that for normal-brightness MSPs a TOA precision of between 80 and 230\\,ns can be achieved at 1.4\\,GHz with 10-minute integrations by the SKA. The actual rms timing residuals for each pulsar will be approximately at the same level only if all the other influences on timing precision (e.g. ISM, spin noise) are either corrected, modelled, or negligible. ",
        "introduction": "Pulsars are stable and rapidly rotating radio sources. This stability \\citep[which rivals the stability of atomic clocks on Earth, see][]{hcmc10} can be exploited through pulsar timing. In brief, pulsar timing works as follows: after the telescope surface focusses the radio signal at the receiver feed, the signal is amplified, sampled and digitised. Subsequently the frequency-dependent dispersion delay caused by the ionised interstellar medium (ISM) is removed in a process called ``de-dispersion''. Then a high signal-to-noise ratio (S/N) pulse profile is obtained by averaging hundreds or thousands of subsequent pulses in a step named ``folding''. Finally this high S/N average profile is matched with an independently obtained standard profile (or an analytic template based on the data itself), in order to derive the time-of-arrival (TOA) of the integrated profile.  Precise monitoring of these TOAs allows the investigation of many aspects of the pulsar, the binary system it inhabits and anything affecting the radio wave propagation. For example, studies of strong gravitational field effects in some binary pulsar systems have in the past enabled some of the most constraining tests of general relativity \\citep{tw89,ksm+06}. Recent progress in software \\citep[e.g.][]{hjl+09} and hardware \\citep[e.g.][]{vbb+10} has led to predictions that within a decade, timing of a group of millisecond pulsars (MSPs) will allow various characteristics of a background of gravitational waves to be determined directly \\citep[e.g.][]{jhv+06,ljp08}. Furthermore, the next generation of radio telescopes (such as the Square Kilometre Array, SKA, and the Five hundred metre Aperture Spherical Telescope, FAST) will have the best chance of detecting the first pulsar--black-hole binary system \\citep[e.g.][]{ckl+04}. Timing of such a system could determine the spin and quadrupole moment of a black hole and then allow direct tests of the Cosmic Censorship Conjecture and the no-hair theorem \\citep{kbc+04}. Pulsar timing at very high precision is required to achieve the aforementioned scientific goals, and it is more readily achieved with MSPs because of their short spin periods and highly stable average pulse shapes. Currently, several MSPs have already been timed at precisions down to a few hundred nanoseconds over time spans of a decade or more \\citep{vbc+09}.  Although the integrated profiles of MSPs appear stable over time scales of years, there are a variety of effects that can affect the shape of an integrated profile on short time scales: multi-path propagation in the turbulent ISM, pulse jitter, data processing artefacts and improper calibration, for example. Profile variations from these effects may only change the pulse shape at low levels, but will cause the subsequent TOA calculation to be less accurate and precise than what is expected if only radiometer noise were contributing to the uncertainty. This will complicate timing with the next generation of radio telescopes since in these cases the timing will be limited by factors other than merely telescope sensitivity.  In order to investigate the level at which short-term instabilities in pulse shape may affect pulsar timing with this new generation of telescopes, we present an analysis on PSR~J0437$-$4715. This pulsar was discovered by \\cite{jlh+93} and is the nearest and brightest MSP known, resulting in outstanding timing precision that has already led to a variety of interesting results \\citep{vbb+01,vbv+08}. Furthermore, the TOA precision of PSR~J0437$-$4715 obtained by current instruments \\citep[see e.g.][]{vbb+10} is already comparable to the precision future telescopes may expect to obtain on other less bright MSPs (see Section \\ref{sec:Conclusions}), making it a perfect target for investigations of the pulsar timing potential of future telescopes.  The structure of this paper is as follows. First we describe the observations and data preprocessing in Section~\\ref{sec:Obs}. Some statistical tools are introduced in Section~\\ref{sec:Tool}. Next we review the possible effects involved in profile distortion and present the results of data reduction in Section~\\ref{sec:Issues}. We conclude with an overview of our main findings, prospects for the precision timing with the next generation of radio telescopes, and a brief discussion of future research in Section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions} \\subsection{Summary}\\label{ssec:sum} In this paper, we have used the brightest and most precisely timed MSP, PSR~J0437$-$4715, to illustrate the most important phenomena known to affect the shape of pulse profiles; to estimate the impact of these phenomena and to evaluate the efficacy of mitigation schemes. By concentrating on the brightest MSP currently known, we are able to cast some light on effects that future telescopes like FAST or the SKA will come across as a matter of course in any standard MSP.  We find that pulse phase jitter may be dominant in current short-term timing of PSR~J0437$-$4715, though further analysis is required to enable quantification of this effect. The effects of faulty de-dispersion are found to be small, negating the need for frequent updates of the dedispersion DM. The effects of scintillation across the bandwidth could not be fully investigated given the limited bandwidth of our data, but usage of frequency-dependent templates can be expected to resolve the aforementioned problem that scintillation might cause. We present the results of the application of correction algorithms for the most important digitisation artefacts and illustrate the importance of full calibration modelling, as opposed to the more traditional calibration for differential gain and phase only. Finally, we point out the importance of adequate monitoring of the digitisation statistics for 2-bit pulsar observing systems, which is not commonly practised.  We further propose a few diagnostic plots to assess the data quality of any pulsar timing data: \\begin{itemize} \\item Time-$\\ss$ curve: Any non-Gaussian variations in profile sharpness suggests changes in pulse shape and imply the data should be carefully studied before being included in any timing analysis (e.g. Fig.~\\ref{tsharp t2bit}). \\item S/N-$\\sigma_{\\rm TOA}$ relation: To assess the quality of the standard profile used (Fig.~\\ref{tmpl matching simulation}). Deviations from the theoretical relationship or significant differences between datasets indicate data distortion (e.g. Fig.~\\ref{snrsigmas}). \\item S/N-$\\sigma_{\\rm TOA}$ for sub-bands: As in the previous point, any mismatch between curves of different sub-bands indicates loss of data quality (e.g. Fig.~\\ref{snrsigmamultifreq}). \\end{itemize} Note that the analysis is applicable not only to future telescopes, but also to existing ones such as Parkes, Arecibo, Effelsberg, and the Large European Array for Pulsars \\citep[LEAP, see][]{fvb+10}, if the system sensitivity is high enough given the brightness of the source.  \\subsection{Timing in the new millennium} Since PSR~J0437$-$4715 is two orders of magnitude brighter than most MSPs, current observations of this pulsar provide a good demonstration of the TOA measurement situation for most of the MSPs by the next generation of radio telescopes. The 21$^{\\rm st}$ century radio telescope will significantly increase the S/N of pulsar detections and correspondingly reduce the uncertainty of TOA measurements through vast increases in effective collecting area.  Fig.~\\ref{sigmasimu} shows the expected MSP TOA precision that can be achieved by SKA and FAST. Here for Parkes and FAST we only consider the uncertainty induced by radiometer noise and pulse jitter at the worst-case level derived in Section~\\ref{sssec:realjitter}, while for the prediction of SKA we calculate the upper limit by applying the worst-case in jitter and set up the lower limit by assuming no pulse jitter at all. The instrumental effects causing profile distortions and phase fluctuations are neglected based on previous analyses and the ISM influence on pulse shape is assumed to be either corrected \\citep[e.g.][]{wksv08} or not significant on a short timescale with provided frequency and bandwidth. Concerning the properties of MSPs, we assume a mean flux density of 3.0\\,mJy at 1.4\\,GHz, 50 MHz bandwidth, spin period of 5.0\\,ms and 5\\% pulse width. In addition, a 10\\,K sky temperature at 1.4\\,GHz is applied. Consequently, it is shown that the TOA precision for normal MSPs can be improved by over one order of magnitude with the next generation of radio telescopes, even when assuming a worst-case level for pulse phase jitter. The result indicates that jitter induced uncertainties will be considerable in future timing of MSPs. The instrumental parameters used for these calculations are presented in Table~\\ref{tab:FutTel} \\citep{sac+07,nan06}.  We conclude that at 1.4\\,GHz for 10-min integrations, a TOA precision of between 80 and 230\\,ns can be expected in timing of normal brightness MSPs by SKA. The wide range of this prediction is caused by our limited knowledge of the pulse jitter mechanism. The future radio telescopes will enable deeper investigation of profile stability of single pulses. Presently the comparatively high levels of radiometer noise mean that only a few bright MSPs have allowed a systematic study of their single pulses, where only a subset of the single pulses can be clearly detected \\citep{cstt96,jak+98,es03}. The significant improvement in instrumental sensitivity will both dramatically increase the number of MSPs available for single pulse analysis and reduce the selection effects present in current studies.  \\begin{figure} \\includegraphics[scale=0.45]{sigmapredic.eps} \\caption{Expected average-brightness MSP TOA scatter versus integration time for different instruments. Only radiometer noise and pulse jitter are accounted for as influences on the measurement accuracy. \\label{sigmasimu}} \\end{figure}  \\begin{table} \\centering \\caption{Key parameters used in prediction of TOA precision for different instruments at 1.4\\,GHz.}\\label{tab:FutTel} \\begin{tabular}[c]{ccc} \\\\ \\hline &$A_{\\rm eff}$ ($\\rm m^{2}$)  &$T_{\\rm sys}$ (K)\\\\ \\hline Parkes  &$2.2\\times10^{3}$       &23.5   \\\\ FAST    &$7.1\\times10^{4}$       &20.0   \\\\ SKA core&$2.5\\times10^{5}$       &30.0   \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{Limitations and future work} The analysis based on PSR~J0437$-$4715 is limited by the selected sample. Note that the DM and scattering timescale (see Section~\\ref{ssec:ISM}) of this pulsar are relatively low so that the influence of the ISM on the TOA precision within the provided bandwidth is expected to be negligible. As system sensitivity and observational bandwidth are improved, more high-DM sources will be timed at high precision TOA measurements. The effects of scattering and scintillation on profile shape will then become more considerable than shown in this paper, and will need to be accounted for by using correction algorithms \\citep{bcc+04,wksv08} and/or frequency-dependent template profiles."
    },
    "1107/1107.2049_arXiv.txt": {
        "abstract": "We present optical ($UBVJ$), ultraviolet (FUV, NUV), and high resolution atomic hydrogen (\\HI) observations of the nearby blue compact dwarf (BCD), VII Zw 403. We find that VII Zw 403 has a relatively high \\HI\\ mass-to-light ratio for a BCD.  The rotation velocity is nominally $10-15$ \\kms, but rises to $\\sim 20$ \\kms\\ after correction for the $\\sim 8-10$ \\kms\\ random motions present in the gas. The velocity field is complex; including a variation in the position angle of the major axis going from the NE to the SW parts of the galaxy. Our high resolution \\HI\\ maps reveal structure in the central gas, including a large, low-density \\HI\\ depression or hole between the southern and northern halves of the galaxy, coincident with an unresolved x-ray source.  Although interactions have been proposed as the triggering mechanism for the vigorous star formation occurring in BCDs, VII Zw 403 does not seem to have been tidally triggered by an external interaction, as we have found no nearby possible perturbers. It also doesn't appear to fall in the set of galaxies that exhibit a strong central mass density concentration, as its optical scale length is large in comparison to similar systems. However, there are some features that are compatible with an accretion event: optical/\\HI\\ axis misalignment, a change in position angle of the kinematic axis, and a complex velocity field. ",
        "introduction": "VII Zw 403 is classified as a blue compact dwarf (BCD) galaxy. Compact galaxies were first identified by \\citet{zwicky64} as high surface brightness objects that were just distinguishable from stars on the Palomar Sky Survey.  This original classification included galaxies with a wide range of colors and total luminosities. Later, \\citet{tm81} presented a study of 115 blue compact dwarf galaxies, which they defined as low luminosity ($M_B \\leq -18$) galaxies with a compact optical size ($\\lesssim 1$ kpc in diameter), that have optical spectra with ``strong sharp narrow emission lines superposed on a blue continuum.'' The last two properties were chosen to ensure that the objects in their study had dense regions with  active star formation.  It was shown early on that BCD galaxies could not have maintained such vigorous star formation for long periods of time \\citep{searle72}, so such activity must be episodic. The cause of this activity is unclear. Unlike spiral galaxies, triggering mechanisms such as spiral density waves and shear are not present in dwarf galaxies, which tend to be slow, mostly solid-body rotators. Gravitational instabilities in the gas, which work well to predict the location of star formation in large galaxies, seem to work differently in dwarfs, where the average gas density is typically below the critical \\citet{kennicutt89} value by a factor of 2 or 3, and yet, stars still form.  \\citet{he04} showed that the nature of the \\HII\\ regions in BCDs are not that different from those in Ims; they are just crowded more closely together towards the center in the BCD galaxies. The amount of star formation caused by pressurized triggering versus spontaneous instabilities is likely similar as well. So what causes the centralized star formation bursts that characterize BCDs?  Two possibilities explored in this paper include an inherent central mass density concentration \\citep{salzer99,vanzee01} or gravitational interactions with nearby objects, including gas infall/accretion \\citep{taylor97, pustilnik01, he04, brosch04}.  Although gravitational interactions are often invoked to explain bursts of star formation in galaxies, previous studies have shown that not all the star formation activity in BCDs can be explained by external gravitational triggers. A plot of current normalized star formation rate vs.\\ distance to the nearest projected (catalogued) galaxy \\citep[see Figure 14 in][] {he04} shows no obvious correlation between presence/distance to companion and star formation activity for Ims, BCDs, or Sms in a sample of 140 galaxies. This is consistent with other searches for bright companions \\citep{campos91,telles95,telles00}, and with work done by \\citet{brosch04}. This does not rule out possible triggering by much fainter companions, though, and there is some evidence for an excess of small/faint companions to BCDs \\citep{taylor97,pustilnik01}, and there are several examples of BCDs with nearby small galaxies or gas clouds \\citep[e.g.][]{putman98,pustilnik03}.  Such faint companions are difficult to detect and may have been missed in earlier studies of BCDs. As pointed out by \\citet{he04}, dwarf galaxies are common so ``the most likely object to perturb any galaxy is a dwarf,'' but the effect of a small/faint companion on a nearby dwarf galaxy will be more dramatic than on a much larger system. They suggest that a possible mechanism for driving mass inflow in dwarfs would be perturbation by a small companion, resulting in angular momentum redistribution, and that such interactions could have happened long ago. Once a low angular momentum system is created, it can continue to host episodic star formation for long time periods \\citep[e.g.][]{vanzee01}.  There is some evidence for central mass density concentration in BCDs in both the neutral gas \\citep{taylor94,taylor95,vanzee98,simpson00} and in the optical underlying host galaxy \\citep{papa96a,papa96b,salzer99}.  \\Citet{vanzee01} find that the 6 BCDS in their sample have steep rotation curves relative to low surface brightness dwarfs (LSBDs) with similar luminosities, with their mass distributions (stars, gas, and dark matter) more centrally concentrated, and that they are lower angular momentum systems than similar LSBDs. They suggest that it is just such low angular momentum systems that are able to collapse into centrally concentrated objects with high central densities. For a solid body rotation curve, the star formation threshold density is proportional to the slope of the rotation curve in a Toomre disk instability analysis \\citep{toomre64}. Hence, a steep rotation curve raises the critical threshold density for star formation. This delays the onset of star formation until the gas densities are quite high; and once star formation does begin, the system is more susceptible to bursting behavior than a more diffuse (higher angular momentum) system.  Another possible cause for the star formation activity that characterizes BCDs is the accretion/merger of a small gas cloud. Such events have been invoked to explain mismatches between the optical and (gas) kinematic axes in galaxies, twists in the velocity field, extended/unusual gas morphology \\citep[e.g.][]{hunter02}, counter-rotation of the stellar and gas components, and the presence of stellar ``shell'' features in gas-poor systems \\citep[e.g.][]{coleman04}. Diagnosing the occurrence of a previous accretion/merger event in dIrr galaxies, however, is complicated by their low gravitational potentials, large gas fractions, and slow rotation.  With slower rotation and low gravitational potentials, their velocity fields are more easily disturbed by the energy injection from star formation regions, for example.  As part of a project to determine what regulates star formation in small galaxies \\citep{he04,he06}, we have obtained multiwavelength data ($UBV$ and \\ha\\ images) for 136 galaxies, with additional data (e.g. $JHK$, NUV, FUV) for a subset.  These data have recently been augmented by deep, high resolution interferometric atomic hydrogen spectral line (\\HI) observations of 40 of these galaxies as part of the LITTLE THINGS (LT) project: Local Irregulars That Trace Luminosity Extremes: The \\HI\\ Nearby Galaxy Survey (http://www.lowell.edu/users/dah/littlethings/).  As part of our investigation into star formation in dwarf galaxies, we are interested in looking at ``what goes wrong'' in BCDs---why they have such strong central star formation occurring. \\vii\\ is one of a few BCD galaxies in our sample, and we use it here as a case study for actively star-forming dwarfs. Because it is apparently relatively isolated, with no spiral density waves or shear, it is a good testbed for examining how star formation might be triggered in active dwarf galaxies. We present the results of an initial study here; future papers associated with the LT project will present more in-depth analysis and modeling, particularly of the kinematics.  Although \\vii\\ has generally been assumed to be a member of the M81 group, \\citet{kara02} state that VII Zw 403, with a distance of 4.4 Mpc, is behind the M81 group and is being accelerated towards both it and the Local Group. They find that VII Zw 403 is 1.14 Mpc from M81 and that the M81 group has a zero-velocity radius of $1.05 \\pm 0.07$ Mpc. \\citet{lynds98} and \\citet{regina99} determined a distance of 4.5 Mpc using the tip of the red giant branch from color-magnitude diagrams based on {\\it Hubble Space Telescope} (HST) data. This places it on the far side of the M81 group. \\citet{he04} used the NASA/IPAC Extragalactic Database (NED)\\footnote{http://nedwww.ipac.caltech.edu/} to search for nearby galaxies in the plane of the sky within 1 Mpc and $\\pm 150$ \\kms. The only object found (an Im galaxy, KDG 073) had a projected distance of 900 kpc with a velocity difference of $-32$ \\kms.  \\citet{loose85} classified \\vii\\ as an iE galaxy: one with irregular bright star formation regions near, but not at the center of, a low-surface brightness envelope. Subsequent observations with HST confirmed the presence of a smooth elliptical distribution of old, metal-poor red giant stars with an exponential fall-off in surface density \\citep[and references therein]{lynds98,regina99}. The scale-length for this evolved population is about 1 kpc \\citep{lynds98}, which is normal for BCDs \\citep{papa96a}. There are several bright \\HII\\ regions with young blue stars loosely clustered nearby, and diffuse \\ha\\ emission as well. Fabry-Perot \\ha\\ maps \\citep{thuan87} detected little rotation, finding only a chaotic velocity field and a velocity dispersion of approximately 30 \\kms\\ in the diffuse emission, consistent with stellar winds.  Initial \\HI\\ inteferometric observations by \\citet{thuan04} revealed that although the \\HI\\ is centrally concentrated in general in the galaxy, there are two areas of higher flux density. The velocity dispersion map showed high values for the region separating the two peaks, and they interpret the high dispersion in the region between the peaks as due to a lack of flux. As most \\HI\\ observations of BCDs show smoothly distributed gas with only one mostly centralized density peak, the presence of two peaks was intriguing, especially if truly separated by a region of lower density. To examine the gas morphology and kinematics in greater detail, and to search for possible companions and/or evidence of gravitational interactions, we subsequently obtained high spatial and velocity resolution \\HI\\ interferometric data of \\vii\\ to combine with multiwavelength optical, UV, and IR data. These data are presented below. ",
        "conclusions": "\\subsection{Relation between Star Formation and Gas}  As can be seen in Figure~\\ref{fig:bcd0m0onha}, the highest \\HI\\ column density region roughly corresponds to the region of strongest \\ha\\ emission. This is seen in other BCDs as well \\citep{vanzee98}. There is diffuse \\ha\\ emission that spreads to the north, but not as far as the \\HI\\ emission. \\citet{thuan04} find that there is weak continuum emission (1.4 GHz) associated with the \\ha\\ emission; with a slight extension to the north at 1.4 GHz as well. Without a spectral index, they were unable to verify that the continuum emission is mainly thermal in nature (and thus directly due to to the \\HII\\ region), but it does seem likely.  \\citet{ott05a,ott05b} presented {\\it Chandra} x-ray observations of several dwarf starburst galaxies, including \\vii. Their data show no diffuse x-ray emission associated with \\vii. They found that the diffuse x-ray emission they detected in the other galaxies in their sample was probably associated with the development of galactic winds from coronal gas heated by star formation.  They suggest that the lack of detected diffuse x-ray emission for \\vii\\ (they give an upper limit of $ 2.8\\times 10^{38}$ erg s$^{-1}$) could be due to several possible factors, including low metallicity or smaller volume.  The expanding \\ha\\ shells in the two galaxies with no detected diffuse x-ray emission (\\vii\\ and I Zw 18) are 10-17 Myr older than those in the galaxies for which they do detect diffuse emission. Perhaps a coronal wind never developed, or perhaps it has already had time to cool.  \\citet{ott05a} did detect one unresolved x-ray source whose power-law spectral index is consistent with that of an x-ray binary, but did not detect the filaments or lobes seen by \\citet{papa94} from ROSAT data. Similar x-ray sources in the \\citet{ott05a} sample of 8 starbursting galaxies are located near ``bright \\HII\\ regions, rims of superbubbles or young stellar clusters.'' The location of the unresolved x-ray source relative to the \\ha\\ emission can be seen in Figure~\\ref{fig:bcd0m0onha}. The x-ray source is located near, but not coincident with, an \\ha\\ knot. As seen in Figures~\\ref{fig:bcd1m2onha} and \\ref{fig:bcd0m0onha}, the x-ray source is also on the north edge of the high dispersion knot, and is located in/on the northeast side of the depression/cavity in the \\HI, in a smaller region of even lower density. Recall that the broadband surface photometry (Figure~\\ref{fig:sb}) was consistent with constant star formation in the outer regions over the past 10--20 Gyr, with more recent star formation (past 200 Myr) in the central regions. If the x-ray source is, indeed, an x-ray binary, then perhaps we are seeing an evolved stellar system located near a region of recent star formation, and also near a region of disturbed neutral gas. The logical explanation is that a star formation episode and/or resulting supernova has injected energy into the \\HI, increasing the dispersion; and depleting or perhaps clearing out some of the neutral gas.  To test this, we examined the kinematics of the gas in the \\mres\\ cube in a small region around the x-ray source. To enhance the signal-to-noise in the spectra, we ``collapsed'' a six-pixel square region centered on the x-ray source to a single pixel, and then examined the resulting spectrum. There was a fairly clear double-peak visible, so we fit a pair of gaussians to them (using {\\sc xgaus} in {\\sc aips}). The two peaks are separated by 13.5 \\kms. This is consistent with expansion (or contraction) of $\\sim 7$ \\kms. The apparent radius of the depression around the x-ray source in the \\mres\\ data is $\\sim 105$ pc, which is roughly the same size as the beam. This means that the region may not be resolved, so this size is an upper limit. If we assume a constant expansion of 7 \\kms, we get approximately 15 Myr to clear out a depression of that size.  We then tried to estimate the amount of energy required to form such a depression using the method in \\citet{chevalier74}. Here we assume a spherically symmetric expansion. If the expansion has already blown out of the disk, this may not be the case (and would increase the estimated age of the feature; see below). For a velocity of 7 \\kms\\ and a radius of 105 pc, we find the total energy required to be $E_s=1.6n_0^{1.12} \\times 10^{51}$ ergs. To get an estimate for the density, $n_0$, we assumed a pre-existing column density of $1.5 \\times 10^{21}$ \\acm\\ based on the average density in the region around the small depression, and a scale height for the galaxy of 200 pc. The scale height was taken to be approximately equal to the radius of the larger \\HI\\ depression/cavity (see below). This gives us $n_0=1.0$ cm$^{-3}$, and a resulting total energy of $1.6 \\times 10^{51}$ ergs. This is similar to the $10^{51}$ ergs from a typical Type II supernova, so it is feasible in this sense.  We then attempted to apply the models of \\citet{mccray87} for the formation of supershells from stellar winds and supernovae to estimate the number of stars with masses greater than 7\\msun\\ required to produce a hole of a given size. The resulting number was approximately 200, of which about 50 should be O stars. However, the resulting timescale for shell formation calculated from the models was short, only 2 Myr. In this short time, the expansion due to the star formation is still primarily from stellar winds; and it is unlikely that any stars would yet have exploded as supernovae, and we should be able to see an OB association of this size---which we don't. So about all we can conclude is that the x-ray source appears to be coincident with an expanding depression in the \\HI; and adjacent to a region of disturbed gas to the south.  We also examined the kinematics of the gas around the edges of the larger depression/cavity in the \\HI. However, no expansion was detected. If we now use the average velocity dispersion in the gas (9 \\kms) as an upper limit to any expansion, with a radius of 205 pc we obtain a lower limit of 22 Myr for a hole this size to form. Because we don't detect expansion, the hole could be stalled because it has broken out and is freely expanding perpendicular to the disk, and so could be older than that. We used the same methods to estimate the required energy and number of stars required to form such a depression as for the area around the x-ray source. This time, we used a column density of $2.5 \\times 10^{21}$ \\acm\\ based on the existing gas density around the cavity; giving us an $n_0$ of 1.6 cm$^{-3}$. The required energy is then $3 \\times 10^{52}$ ergs (equivalent to approximately 30 SNe). However, from the \\citet{mccray87} equations, the estimated number of stars needed to create such a large depression via winds (and possibly supernovae, but see the timescale discussion, below) with masses greater than 7\\msun\\ is then 2800, which is extremely high. If there were a young cluster that large, it would be easily visible. This suggests an older age for the formation of the large \\HI\\ cavity; or perhaps the gas was slowly depleted/blown-out over long time-scales by a series of star-forming events rather than by one large cluster.  A more serious problem with applying the \\citet{mccray87} model is that doing so results in a calculated timescale of only 3 Myr; which is not only  short in stellar evolution terms, but also inconsistent with the lower limit of 22 Myr we calculated, above. A similar disconnect was found by \\citet{lynds98} for formation of the large \\HII\\ regions in \\vii. They used the Wide Field Planetary Camera 2 (WFPC2) on the {\\it Hubble Space Telescope} (HST) to produce color-magnitude diagrams (CMDs) of several of the bright \\HII\\ regions in \\vii.  The brightest \\HII\\ region (visible as the darkest region of the grayscale in Figure~\\ref{fig:bcd0m0onha} and marked in Figure~\\ref{fig:bcd-1m0onhst}) was consistent with an OB association containing approximately 25 bright blue (O type) stars; the age of the cluster based on the CMD was 4--5 Myr. From their ground-based long-slit emission line spectra of the \\HII\\ regions, they were able to detect expanding \\ha\\ shells; including a partial shell around the brightest \\HII\\ region with a diameter of 79 pc and an expansion velocity of $\\sim 65$ \\kms. They then attempted to use the \\citet{mccray87} models as we did, but they also found a timescale that was too short: $\\sim$ 1 Myr, which did not match the 4--5 Myr age estimate from the CMD. They concluded that perhaps this indicated that the starburst that formed the association was not instantaneous \\citep{shull95}.  \\citet{lynds98} also produced CMDs of some of the central areas of \\vii. The location of their ``center'' and ``north'' fields are marked in Figure~\\ref{fig:bcd-1m0onhst}, which shows the HST image with contours from the \\hres\\ integrated flux map. Neither field corresponds exactly with the \\HI\\ cavity/depression, but the ``north'' field is located on the east side of it, and just south of the x-ray source. The CMD for the ``north'' field is less rich in massive stars than the one for the ``center,'' which revealed approximately 20 red supergiants ($> 10$ Myr) and some blue supergiants ($\\sim 5$ Myr), but both regions contain a mixed-age population of stars with ages of 5 and 10 Myr. In general, \\citet{lynds98} find a relatively smooth elliptical distribution of the evolved stars with a stellar surface density that falls off exponentially. The younger, bluer massive stars tend to exhibit clustering into associations and groups, and are primarily located in the central regions of the galaxy. The youngest are associated with the bright \\HII\\ regions, which we find are co-located with \\HI\\ peaks (Figure~\\ref{fig:bcd-1m0onha}). They posit star formation 1--2 Gyr ago followed by a strong starburst episode 600--800 Myr ago, with less vigorous star-formation since then. Interestingly, the kinematic center of \\vii\\ as identified by the rotation curve analysis lies  close to the center of the large \\HI\\ cavity. If the cavity is the result of star-formation, this would be consistent with centrally-concentrated star-formation activity in the past. There is evidence that stellar feedback from star formation events can and does create sometimes large ($\\approx$ 1 kpc) \\HI\\ holes and shells, as found by \\citet{weisz09} in IC 2574.  Another possibility for forming the hole is consumption of the gas by star formation. The timescale to consume the gas that filled the hole at the current star formation rate, given our estimate of the pre-existing level of the gas in that region, is of order 700 Myr. This is relatively short if star formation has been on-going in this region for that extended period of time. Alternatively, the original gas in this region could have been consumed quite quickly by the starburst that took place 600--800 Myr ago \\citep{lynds98}, but we then require that the hole not be replenished very rapidly, which seems unreasonable given that the hole should have filled in long ago just from the dispersion in the gas.  \\subsection{Origin of the Starburst}  One possible explanation for the enhanced star formation episode that characterizes BCDs is proposed by \\citet{salzer99}, who suggest that it is a natural consequence of the central mass density concentration found in BCDs. In an optical study of a sample of 18 BCDs and 11 dIrrs, they find that when the luminosity contribution from the starburst is ignored (by fitting the underlying host galaxy's surface brightness profile as determined from an exponential fit to only the outer portions), the underlying hosts of BCDs have shorter scale lengths and higher central surface brightnesses than dIrrs with similar absolute magnitudes. They cite the shorter scale lengths as evidence of stellar compactness: that BCDs are low-surface brightness galaxies with the highest central stellar mass concentrations in the dwarf spectrum.  There is some evidence that this may be the case for the \\HI\\ gas in BCDs as well: \\citet{vanzee98} plotted \\HI\\ surface density profiles for a sample of BCDs and dIrrs, and found that the gas in BCDs is more centrally concentrated than in similar dIrrs. It is commonly thought that the gas surface density must exceed a critical threshold in order to support a starburst\\footnote{Although azimuthal averages of the gas density in dwarfs rarely exceed this critical value, there is evidence that local values in the region of active star formation do \\citep{hew01,vanderhulst93,vanzee97,meurer98,deblok06}.} \\citep{kennicutt89}; \\citet{salzer99} also suggest that the high central mass density concentration allows these galaxies to retain their gas, allowing episodic bursts, and calculations by \\citet{maclow99} indicate that mass loss should be minimal for galaxies with gas masses greater than $10^7$ \\msun. A follow-up study of the \\HI\\ in a sample of BCDs by \\citet{salzer02} found that the $M_{HI}/L_B$ for these galaxies is about twice as high as that for low surface brightness dwarfs  when only the luminosity of the underlying low-surface brightness component of the BCD is taken into account. This provides a significant gas reservoir to support the starburst episodes, and could enhance the ability of the galaxy to retain gas during/after a burst. Additionally, if the gas remains concentrated even after a burst, then the timescale between bursts could be  short. This would make it difficult to find a BCD in the ``off'' state.  If this hypothesis is correct then we would expect \\vii\\ to exhibit signs of central compactness, both optically and in the \\HI. We find that compared to the sample in \\citet{salzer99}, \\vii\\ has a slightly lower $\\mu_{B_0}$ at a given M$_B$ than their sample of BCDs, although slightly higher than for their LSBDs. As mentioned in \\S\\ref{sec:optical}, \\vii's scale length (relative to its $M_V$) lies at the peak of the distribution of the sample of dIms and BCDs plotted in \\citet{he06}. Compared to the \\citet{salzer99} sample, however, it is large not only for their sample of BCDs, but also compared to their LSBD sample. So the underlying stellar component does not appear to be greatly centrally concentrated. However, there are some indications for central concentration in the \\HI\\ on larger scales (i.e.\\ at a linear resolution similar to that for which structure is seen in the \\HI\\ in dIrrs; see Figure~\\ref{fig:cd1m0onv}).  However, since the vigorous star formation episodes that characterize BCDs are relatively short-lived compared to the age of the host galaxy, perhaps these systems don't start out with a high central mass concentration. \\Citet{vanzee01} finds that BCDs are low angular momentum systems, and suggests that they have collapsed to more centrally concentrated systems with steeper rotation curves. This affects the critical threshold density for star formation, and makes BCDs more susceptible to star bursts than higher angular momentum systems. This ``collapse'' requires a redistribution of mass in the galaxy.  One proposed mechanism for mass redistribution and angular momentum transfer in low mass BCDs (without nearby companions) is a barlike torque arising from the irregular symmetry and irregular mass distribution in these small systems, or from an elongated dark matter halo \\citep{he04}. Such mass flow should show up in the rotation pattern of the gas disk as a central twisting of the isovelocity contours. \\vii\\ has a change of 15--20\\arcdeg\\ between the NE and SW parts of the galaxy, so we can't rule out the possibility of some kind of mass flow in the gas disk.  Another possible mechanism for flattening the gravitational potential and expansion of the stellar component is suggested by \\citet{papa96b}, who propose that the heating of the ISM by the energy from the starburst could be the mechanism for mass and angular momentum redistribution. Although \\vii's scale length is not particularly short, which could indicate that it has undergone some expansion, the lack of diffuse x-ray as well as the lack of extensive large-scale structure in the \\HI\\ (the \\HI\\ depression discussed above occupies only a small, central part of the galaxy) would seem to indicate that not much energy has been injected into the galaxy as a whole by star formation. It seems unlikely, then, that enough energy would have been produced to initiate any significant redistribution of the mass in the galaxy, especially the stellar component. In conclusion, it isn't clear that this theory of large-scale mass re-arrangement can explain what we see in \\vii.  Interactions, either with other galaxies or with small gas companions, have also been proposed as the triggering mechanism for the starbursts that occur in BCDs (see \\citealt{brosch04} for an overview of the literature), yet surveys to detect nearby possible perturbers around BCDs have had mixed results \\citep[e.g.][]{campos91,telles95,taylor97,putman98,telles00,pustilnik01,pustilnik03,brosch04,he04}, \\vii\\ would seem to be in the set of galaxies that haven't been tidally triggered, as we have found no nearby \\HI\\ clouds or companion galaxies. However, if it had accreted a small perturber in the distant past, this might be difficult to detect.  There is observational evidence that some dwarf systems have undergone accretion or merger events in the past. A nearby example is thought to be IC 10, which \\cite{wilcots98} argue still has primordial gas settling into the galaxy via accretion. Classic signatures of accretion include the appearance of two distinctly different kinematic components, visible in the \\HI\\ or stellar velocity fields. Although the gas isovelocity contours for VII Zw 403 are somewhat perturbed, this isn't unusual for BCDs (\\S \\ref{sec:kinematics}), and there don't appear to be two completely separate components despite the change in position angle of the \\HI\\ kinematic axis. Another accretion signature is a misalignment between the optical and \\HI\\ axes. \\citet{hunter02} plotted the difference between optical and \\HI\\ kinematic axes for 47 Im and Sm galaxies taken from \\citet{swaters99} and found that about half had offsets less than 10\\arcdeg, with the maximum being 40\\arcdeg. For VII Zw 403, the optical axis is at $-11$\\arcdeg. In the NE half of the galaxy, the \\HI\\ kinematical axis is at +45\\arcdeg\\ to +50\\arcdeg, and is around +20\\arcdeg\\ to +30\\arcdeg\\ in the SW part, so there is a somewhat significant misalignment between the gas and optical components of between 30\\arcdeg\\ and 60\\arcdeg.  For further comparison, we look to the gas-rich Im galaxy DDO 26, for which \\citet{hunter02} conclude that the most likely explanation for the lack of relaxation they detect in its gas kinematics is the on-going merger of two small gas-rich objects. They find an optical/\\HI\\ axis misalignment of 60\\arcdeg, along with a change in the \\HI\\ position angle with radius. They also find that the velocity field of DDO 26 is fairly regular in the inner regions and becomes less so in the outer disk. These are features that are similar to but more extreme than those we see in \\vii. Differences include the presence of a faint \\HI\\ arm in DDO 26 and the detection of two separate peaks in the \\HI\\ intensity which were both fairly broad ($\\sim 40$ \\kms) in velocity. Additionally, DDO 26 is currently quiescent, with a SFR of 0.0004 \\msun\\ yr$^{-1}$ kpc$^{-2}$ compared to \\vii's rate of 0.02 \\msun\\ yr$^{-1}$ kpc$^{-2}$. So from these comparisons, it seems possible that \\vii\\ underwent an accretion or merger event in the past.  Another possibility to explain the skew in the \\HI\\ kinematic axis is that expansion of the \\HII\\ regions has caused a subsequent expansion in the surrounding \\HI, lifting it out of the plane of the galaxy. If the near side of the galaxy is to the west (right), and we envision the expanding \\HI\\ as distributed in a torus aligned with the plane of the galaxy, then gas to the left of the kinematic center would be blue-shifted relative to the bulk of the galaxy. Likewise, the gas to the right of center would be somewhat behind the bulk of the galaxy and red-shifted. The expanding gas would then skew the observed velocity field in the sense that we see: the isovelocity contours in the north part will be skewed to the left, and those in the south skewed to the right.  If the expansion is on the order of the 7 \\kms\\ we measured for the lowest density region around the x-ray source, or the 9 \\kms\\ limit set by the average dispersion, then in this slowly rotating galaxy, the expected skew would be large, which is what we see. In this case, the hole in the \\HI\\ would be due to a combination of consumption and slow blow-out."
    },
    "1107/1107.5496_arXiv.txt": {
        "abstract": "{The Milky Way bulge is the nearest galactic bulge and the best laboratory for studies of stellar populations in spheroids based on individual stellar abundances and kinematics. The observed properties point to a very complex nature, which is hard to extrapolate from a few fields.} {We present a method to obtain reddening maps and to trace structure and metallicity gradients of the bulge using data from the recently started ESO public survey Vista Variables in the Via Lactea (VVV). The method is used to derive properties of the fields along the minor axis.} {We derive the mean $J-K_s$ color of the red clump (RC) giants in 1835 subfields in the Bulge region with $-8^\\circ<b<-0.4^\\circ$ and $0.2^\\circ<l<1.7^\\circ$, and compare it to the color of RC stars in Baade's Window for which we adopt $E(B-V)=0.55$. This allows us to derive the reddening map on a small enough scale to minimize the problems arising from differential extinction. The dereddened magnitudes are then used to build the bulge luminosity function in regions of $\\sim0.4^\\circ \\times 0.4^\\circ$ to obtain the mean RC magnitudes. These are used as distance indicator in order to trace the bulge structure. Finally, for each subfield the derived distance and extinction values have been used to obtain photometric metallicities through interpolation of red giant branch colors on a set of empirical ridge lines. The photometric metallicity distributions are  compared to metallicity distributions obtained from high resolution spectroscopy in the same regions.} {The reddening determination is sensitive to small scale variations which are clearly visible in our maps. Our results are in agreement within the errors with literature values based on different methods, although our maps have much higher resolution and more complete coverage. The luminosity function clearly shows the double RC recently discovered in 2MASS and OGLE III datasets, hence allowing to trace the X-shape morphology of the bulge. Finally, the mean of the derived photometric metallicity distributions are in remarkable agreement with those obtained from spectroscopy.} {The VVV survey presents a unique tool to map the bulge properties by means of the consistent method presented here. The remarkable agreement between our results and those presented in literature for the minor axis allows us to safely extend our method to the whole region covered by the survey.}  \\authorrunning{Gonzalez et al.} \\titlerunning{Reddening and metallicity maps of the Milky Way Bulge} ",
        "introduction": "Several studies have attempted to address the problem of the origin of the Galactic bulge, trying to reconstruct its star formation history through the analysis of both stellar ages and chemical abundances \\citep[][]{mcwilliam94,zoccali03,fulbright07,rich_origlia07,lecureur07,zoccali08,clarkson08,brown10,johnson11,gonzalez11}. Early studies were restricted to a low extinction window at a latitude $b=-4^\\circ$ along the minor axis of the bulge, the so-called Baade's Window. However several recent papers presented the evidence for a high complexity of the bulge structure and stellar population content: i) The Milky Way bulge has a radial metallicity gradient \\citep[][]{zoccali08,johnson11} most likely flattening in the inner regions \\citep[][]{rich_origlia07}; ii) it seems to host two kinematically and chemically distinct components \\citep[][]{babusiaux10,gonzalez11}; and iii) it has an X-shape \\citep[][]{mcwilliam10,nataf10}. All these observations strongly argue against the assumption that the properties of the Galactic bulge stellar population in a few low extinction windows could be extrapolated to the whole bulge. Further analysis of the chemical properties of stars in other bulge regions, away from the minor axis, has been partially obtained from the latest sample of microlensed dwarfs from \\citet[][]{bensby10}. Although microlensing provides chemical abundances all across the bulge, a larger sample is required in order to verify the presence and extent of gradients. Consequently, at present, a different approach based on a larger coverage with a statistically significant stellar sample is the key to disentangle the general properties of the bulge.  In this context, a base for such larger coverage studies is required in order to analyze and correctly interpret results. In particular, \\citet[][]{zoccali03} and later \\citet[][]{johnson11} showed how the upper red giant branch photometry can be used to obtain photometric metallicity distributions for a large number of bulge stars in two fields at $b=-6^\\circ$ and $b=-8^\\circ$, respectively. Using a similar photometric approach to derive the metallicity distribution of a much larger bulge region will provide for the first time a general picture of the iron content of the galactic bulge. Unfortunately, in order to reach this goal a first challenge arises from the extinction variations across different bulge sightlines. While most of the studies relied on the large scale extinction maps of \\citet[][]{schlegel}, it is well known that these maps become unreliable when inner regions of the bulge closer to the galactic plane ($|b|<3$) are considered. \\citet[][]{schultheis99}, based on DENIS photometry, and \\citet[][]{dutra03}, based on 2MASS, used the red giant branch photometric properties to obtain 2D extinction maps in several inner bulge regions. \\citet[][]{marshall06} presented 3D models of the extinction properties of the inner Galaxy with a resolution of 15'. However, at this point a homogeneous extinction map of the bulge including both the outer and the most inner regions, where high resolution is required to avoid differential extinction, is still necessary.  Regarding the structure of the bulge, the work of \\citet[][]{nataf10} and \\citet[][]{mcwilliam10}, based on OGLE and 2MASS photometry respectively, provided hints for a X-shape morphology in the outer regions of the bulge which however, seems to merge into a single structure in the innermost regions, i.e. the bar \\citep[][]{rattenbury07}. Both \\citet{mcwilliam10} and \\citet{nataf10} used the observed properties of the RC stars at different lines of sight, which unfortunately restricted the studies to the outer bulge regions due to the non-complete optical coverage of OGLE and the relatively bright limit magnitude of 2MASS. For a thorough study of the bulge structure, and to homogeneously trace both the inner and outer regions, high spatial coverage and deep near-IR photometry are necessary.  Both reddening and structural properties can be obtained from the photometric analysis of RC giants across the Galactic bulge. RC giants are core-helium burning stars, equivalent of the horizontal branch stars in old metal-poor populations, and are easily identified in bulge color-magnitude diagrams. Their intrinsic magnitude has a fairly well understood dependence on population properties such as age or metallicity \\citep{girardi+salaris01}, and reddening effects in color can be clearly identified. For old and relatively metal-rich population such as found in the Galactic bulge \\citep[e.g.][and references therein]{brown10} the population corrections for red clump absolute K-band magnitude are of the order of 0.1~mag \\citep{salaris+girardi02}. For this reason, photometric properties of red clump giants can be used to trace simultaneously extinction and relative distances in the bulge. The VISTA Variables in the Via Lactea (VVV) Survey, mapping 520 deg$^2$ including the southern disk and bulge of the Milky Way \\citep[][]{vvv10}, provides for the first time deep enough near-IR photometry to trace RC stars in all bulge regions including highly reddened innermost parts.  In this paper, we present the first analysis of the Galactic bulge RC giants based on VVV data for a region along the minor axis. We analyze the RC properties in the color magnitude diagrams along different lines of sight to infer reddening properties and mean distances. This information is then coupled with 2MASS photometry of the upper red giant branch (RGB) stars in order to further derive bulge metallicity distributions. This paper describes in detail the method and compares the results with the present literature. In a forthcoming paper we aim to apply the method to the complete 300 deg$^2$ coverage of the VVV bulge area, deriving homogeneous extinction, morphological structure and metallicity maps of the entire Galactic bulge. ",
        "conclusions": "We used the VVV survey data to investigate the properties of the Galactic bulge along the minor axis. Within this bulge region there have been several spectroscopic studies which we use to validate our method. The method is based on the use of properties of the red clump giants. The RC color is used to derive the reddening map along the minor axis extending from $-8^\\circ <b<-0.5^\\circ $. These regions suffer from high extinction and variations on very small scales are observed. Therefore a consistent bulge reddening map obtained from a homogenous tracer all across the bulge, from the innermost crowded and extinct regions to the outer parts is very valuable. Our results for the extinction also show good agreement with maps published by \\citet{dutra03}, and \\citet{schultheis99}.  We have used the de-reddened magnitudes to build the Ks band luminosity functions and to trace the RC along the minor axis. When using the RC as a distance indicator, we find excellent agreement with previous findings of a double red clump for latitudes larger than $b<-5^\\circ$ which trace the X shape of the Galactic bulge \\citep[][]{mcwilliam10,nataf10}. We also detect the additional bump in the luminosity function interpreted as the RGBb by \\citet[][]{nataf11}. This study will  be extended to the complete 300 deg$^2$ bulge region and will provide a consistent trace of the structure of the bulge including the inner regions.  Finally, we used the RC distances and reddening map to build the color magnitude diagrams in the absolute plane and then to obtain individual photometric metallicities for RGB stars in the bulge through interpolation on the RGB ridge lines of cluster sample from \\citet[][]{valenti04}. The final metallicity distributions were compared to those obtained in the spectroscopic studies of \\citet[][]{zoccali08} and \\citet[][]{johnson11} showing a remarkable agreement. We demonstrated clearly the possibility to track the observed metallicity gradient along the minor axis. This provides, within the errors of the method, the opportunity to obtain clues for metallicity gradients along other regions of the bulge. Reddening and metallicity maps for the VVV coverage will be released electronically in the survey website."
    },
    "1107/1107.5669_arXiv.txt": {
        "abstract": "We study the possibility of reheating the universe through the evaporation of Primordial Black Holes created at the end of inflation. This is shown to allow for the unification of inflation with dark matter or dark energy, or both, under the dynamics of a single scalar field. We determine the necessary conditions to recover the standard Big Bang by the time of nucleosynthesis after reheating through black holes. ",
        "introduction": "} The quest for the unification of inflation, dark matter, and dark energy, in different combinations, by a single field, has been studied in Refs.~\\cite{Peebles:1998qn,Lidsey:2001nj,Padmanabhan:2002sh,Scherrer:2004au,Matos:2005yt,Arbey:2006it,Cardenas:2006py,*Cardenas:2007xh,*Panotopoulos:2007ri,Liddle:2006qz,*Liddle:2008bm,Lin:2009ta,Henriques:2009hq,Bose:2009kc,delaMacorra:2009bi,*delaMacorra:2009yi,BasteroGil:2009eb,Lerner:2009xg,*Okada:2010jd,DeSantiago:2011qb}. The main motivation behind all proposals is that we do not yet understand the nature of the components responsible for the three phenomena, but we do know that their special properties are beyond the realm of the ordinary matter described by the standard model of particle physics.  An extreme, most economical, possibility is that all three phenomena can be explained by the existence of one single field. As was first put forward in Ref.~\\cite{Liddle:2006qz,*Liddle:2008bm}, the simplest option at hand is a scalar field $\\phi$ with a potential of the form $V(\\phi) = V_0 + (1/2) m^2 \\phi^2$. The energy scale $V_0$ is to be set at the tiny value of the observed cosmological constant considered in the concordance $\\Lambda$CDM model, and the mass scale of the field, $m \\simeq 10^{-6} m_{\\rm pl}$, is determined by the amplitude of primordial perturbations generated during inflation.  A crucial ingredient in unification schemes that include the inflaton field is a proper and smooth transition from an inflationary era to the Hot Big Bang phase we are all familiar with. There needs to be a \\emph{reheating} phase in which the inflaton field can transfer most of its energy to other relativistic particles, yet be able to survive the reheating process and attain the right amplitude to later become a subdominant {dark matter/dark energy} component in the radiation-dominated era.  The viability of some phenomenological models of reheating for unification models was explored in Refs.~\\cite{Lidsey:2001nj,Liddle:2006qz,Liddle:2008bm}. Some other, more detailed proposals have been explored in the literature, from extra periods of inflation at low energy scales~\\cite{Liddle:2008bm}, to the more popular gravitational particle production at the end of inflation~\\cite[e.g.][]{Ford:1977dj,Spokoiny:1993kt,Peebles:1998qn,Huey:2001ae,*Dimopoulos:2001ix,Bose:2009kc}. This reheating mechanism is largely ineffective, however, because the relative proportion of radiation to scalar field  $\\rho_{\\rm rel}/ \\rho_{\\phi} $ evolves at the same rate as the proportion of gravitational waves to scalar field $\\rho_{\\rm GW}/\\rho_{\\phi}$~\\cite{Sahni:2001qp}. In the case of steep potentials, inflation is followed by a kination era (a possibility first explored in the context of electroweak baryogenesis \\cite{Joyce:1996cp}). If the kinetic energy of the scalar field dominates for some time, the relative density of gravitational waves created during inflation can grow significantly and violate the observational bounds. Specifically, the process of Big Bang Nucleosynthesis (BBN) imposes a condition on the energy density of gravitational waves $\\rho_{\\rm GW}$, BBN requiring that~\\cite{Shvartsman:1969mm,*Carr:1980ab,*Cyburt:2004yc,*Carr:2009jm} \\begin{equation} \\label{gw:condition} \\rho_{\\rm  GW}(t_{\\rm BBN}) < 0.2 \\, \\rho_{\\rm rel}(t_{\\rm BBN}) \\, . \\end{equation} In fact, the BBN bound in Eq.~(\\ref{gw:condition}) has been used to constrain and exclude some models of braneworld inflation and other unification models~\\cite{Giovannini:1998bp,*Boyle:2007zx}.  In this paper, we explore an alternative mechanism to reheat the Universe, exploiting the evaporation of primordial black holes (PBHs) produced at the end of the inflationary phase~\\cite{Barrow:1990he,GarciaBellido:1996qt}, and study its usefulness in unification models. A small initial population of tiny black holes will grow significantly (in relative energy density) in a universe dominated by a stiff fluid. The evaporation of such black holes at the right time would produce the required amount of radiation more effectively than, for instance, the process of quantum particle production in an expanding space-time.  In the PBH reheating scenario, the black hole component effectively reduces the relative density of gravitational waves as \\begin{equation} \\label{rhogw:rhopbhs} \\frac{\\rho_{\\rm GW}}{\\rho_{\\rm PBH}} \\propto \\left(\\frac{a}{a_{\\rm end}}\\right)^{-1}. \\end{equation} Reheating takes place when the energy density $\\rho_{\\rm PBH}$ is transformed into relativistic particles. Therefore, even if the densities $\\rho_{\\rm GW}$ and $\\rho_{\\rm PBH} $ were equivalent at the end of inflation, we would only require 3 $e$-folds of kinetic energy domination to meet the nucleosynthesis requirements. PBH reheating therefore solves the problem described above. Another advantage of reheating through the evaporation of PBHs is that the unification field does not need to be coupled to other matter fields, hence it can later behave as a totally uncoupled dark matter component. PBH reheating for unification purposes was first considered in Ref.~\\cite{Lidsey:2001nj}; see also Ref.~\\cite{BasteroGil:2009eb}.  Our aim here is to determine the general conditions under which {the unification of inflation with either, or both, of the dark fields} can achieve the appropriate reheating with the aid of PBHs and set up the required conditions before BBN.  The featured model of PBH reheating is described in detail in Sec.~\\ref{sec:pbh-reheating}. Together with this description, we establish the conditions that the scalar field potential must fulfil in order to undergo a successful PBH reheating. A particular example of our proposal is illustrated in Sec.~\\ref{sec:examples}. Finally, in Sec.~\\ref{sec:remarks}, we frame our proposal in the context of unification models and draw some concluding remarks. ",
        "conclusions": ""
    },
    "1107/1107.5343_arXiv.txt": {
        "abstract": "We analyze the \\hcop\\ 3-2 and \\hcopi\\ 3-2 line profiles of 27 high-mass star-forming regions to identify asymmetries that are suggestive of mass inflow. Three quantitative measures of line asymmetry are used to indicate whether a line profile is blue, red or neither --- the ratio of the temperature of the blue and red peaks, the line skew and the dimensionless parameter $\\delta v$. We find nine \\hcop\\ 3-2 line profiles with a significant blue asymmetry and four with significant red asymmetric profiles. Comparing our \\hcop\\ 3-2 results to HCN 3-2 observations from Wu et al. (2003, 2010), we find that eight of the blue and three of red have profiles with the same asymmetry in HCN. The eight sources with blue asymmetries in both tracers are considered strong candidates for inflow. Quantitative measures of the asymmetry (e.g. $\\delta v$) tend to be larger for HCN. This, combined with possible \\hcop\\ abundance enhancements in outflows, suggests that HCN may be a better tracer of inflow. Understanding the behavior of common molecular tracers like \\hcop\\ in clumps of different masses is important for properly analyzing the line profiles seen in a sample of sources representing a broad range of clump masses. Such studies will soon be possible with the large number of sources with possible self-absorption seen in spectroscopic follow-up observations of clumps identified in the Bolocam Galactic Plane Survey. ",
        "introduction": "Isolated, low-mass star formation is generally accepted to result from the gravitational collapse of a dense core within a molecular cloud (e.g. Shu, Adams \\& Lizano 1987). In this picture, cores grow by inside-out collapse with material inflowing with higher velocities near the core. Evidence of collapse is observable in self-absorbed, optically thick line profiles. Emission from the infalling envelope on the far side of the protostar will be blueshifted by an amount proportional to the velocity gradient towards the core. This blueshifted emission will not be absorbed by foreground layers that are warmer or at a substantially different velocity (see Walker et al. 1986; Walker, Narayanan \\& Boss 1994), leading to excess emission in the line profile blueward of the source velocity. For moderately optically thick sources, the line will appear skewed blueward while strongly self-absorbed sources will show two distinct peaks with the blue peak moderately to significantly brighter than the red. Large temperature gradients in the core will increase the depth of the self-absorption feature while large velocity gradients will increase the asymmetry of the line.  Before a source with a blue asymmetric line profile can be considered a collapse candidate, other mechanisms for creating a blue profile must be ruled out. Spatially coincident cores that are not physically associated, rotation and outflows can all produce line asymmetries. Emission from an optically thin isotopologue will peak in the middle of the absorption dip of the optically thick line if the emission is from a single source. Rotation and outflows also generate double-peaked line profiles, though their contribution to the line profile is more apparent with maps of sufficient spatial resolution. Even in the absence of such maps, both rotation and outflows should produce equal numbers of red and blue asymmetric line profiles. Because line profiles may simultaneously include contributions from outflows, expansion and rotation in addition to inflow, there are conflicting views on what constitutes unambiguous evidence for inflow (i.e. Snell \\& Loren 1977; Leung \\& Brown 1977). High resolution observations have spatially resolved inflowing, outflowing and rotating components for a small number of sources (e.g. Remijan \\& Hollis 2006, Keto \\& Wood 2006). Other sources are considered strong candidates for infall based on observations of blue asymmetries in multiple molecular lines, probing a range of densities. For a sample of sources, a statistically significant excess of blue profiles suggests that inflow may be an important physical process in determining the line profile.   Previous work on infall has either focused on detailed study of a single source (e.g. IRAS 16293-2422, Walker et al. 1986; Menten et al. 1987; Walker et al. 1988; Narayanan, Walker \\& Buckley 1998; Ceccarelli et al. 2000; Chandler et al. 2005; Remijan \\& Hollis 2006) or statistical studies of a large sample of sources (e.g. Mardones 2003). While multiple statistical studies have examined candidates for infall among low-mass stars (e.g. Mardones et al. 1997; Gregersen et al. 1997; Gregersen et al. 2000; Lee \\& Myers 2011), the literature on higher mass candidates is less extensive and typically only one kinematic tracer is analyzed (e.g. Wu \\& Evans 2003; Fuller et al. 2005; Klaassen \\& Wilson 2008; Sun \\& Gao 2009; Cyganowski et al. 2009). Massive stars form in complex environments and at higher median distances than low mass sources, thus features of individual cores embedded in the larger clump are averaged together in the single-dish beam, making infall motions harder to observe. Observations of high-mass clumps are also complicated by strong temperature gradients and outflows from the multiple cores likely embedded in the clump. However, gas dynamics of the clump envelope may help discriminate between proposed modes of high mass star formation. For example, cores of moderate mass may grow to become massive stars via the global inflow of clump material because of their location deep in to the cluster potential (Smith, Longmore \\& Bonnell 2009).   In this work, we analyze the line profiles of 27 high-mass star-forming regions observed in \\hcop\\ and \\hcopi\\ 3-2 for evidence of inflow. New \\hcopi\\ 3-2 observations are described in \\S2. We examine these line profiles for asymmetry and calculate the excess of blue profiles relative to red profiles in \\S3. Noting that every source with a blue asymmetry in \\hcop\\ 3-2 is also blue asymmetric in HCN 3-2 (Wu \\& Evans 2003; Wu et al. 2010), we compare our \\hcop\\ results to those found for the same sources in HCN 3-2 in \\S4. We discuss these conclusions and their implications for the advent of ALMA in \\S5. ",
        "conclusions": "Unlike the low mass sources studied by Gregersen et al. (1997) and Mardones et al. (1997), these sources are high mass clumps that likely have multiple cores embedded in the larger cloud. Adopting the language of Wu \\& Evans (2003), we argue that the eight sources displaying a blue asymmetric profile in both HCN and \\hcop\\ are candidates for inflow in order to make clear that this is not evidence of the gravitational collapse expected for a single core (Shu 1977). Single dish observations in two tracers are not sufficient to confirm inflow, however, the eight sources that show a blue asymmetry in both lines are strong inflow candidates and merit follow-up observations to test this hypothesis. Higher resolution maps will be useful for disentangling contributions to the observed line profile as blue asymmetries due to infall should be strongest towards the center of a source (a single core) while blue peaks due to outflows may peak some distance away from the continuum peaks (Mardones et al. 1997).  Ultimately, identifying the best tracers and candidates for inflow in high-mass regions lays the ground work for higher resolution studies with instruments like ALMA that will be able to resolve individual clumps. Hundreds of sources with asymmetric (and possibly self-absorbed) \\hcop\\ 3-2 line profiles have been uncovered in molecular line follow up of sources discovered in the Bolocam Galactic Plane Survey (Schlingman et al. 2011, in press, BGPS; Aguirre et al. 2011). Of the 1882 sources already observed in \\hcop, $\\sim11$\\% show possible self-absorption. The full sample of $\\sim6000$ sources should provide hundreds of inflow candidates for high spatial resolution follow-up."
    },
    "1107/1107.0616_arXiv.txt": {
        "abstract": "% { It has become increasingly apparent that studying how dark matter haloes are populated by galaxies can provide new insights into galaxy formation and evolution. In this paper, we present a detailed investigation of the changing relationship between galaxies and the dark matter haloes they inhabit from $z \\sim 1.2$ to the present day. We do this by comparing precise galaxy clustering measurements over 133~$\\deg^2$ of the ``Wide'' component of the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS) with predictions of an analytic halo occupation distribution (HOD) model where the number of galaxies in each halo depends only on the halo mass. Starting from a parent catalogue of $\\sim3\\times10^6$ galaxies at $i'_{\\rm AB}<22.5$ we use accurate photometric redshifts calibrated using $\\sim10^4$ spectroscopic redshifts to create a series of type-selected volume-limited samples covering $0.2 < z < 1.2$.  Our principal result, based on clustering measurements in these samples, is a robust determination of the luminosity-to-halo mass ratio and its dependence on redshift and galaxy type. For the full sample, this reaches a peak at low redshifts of $M^{\\rm peak}_{\\rm h}=4.5\\times10^{11} h^{-1} M_\\odot$ and moves towards higher halo masses at higher redshifts. For redder galaxies the peak is at higher halo masses and does not evolve significantly over the entire redshift range of our survey. We also consider the evolution of bias, average halo mass and the fraction of satellites as a function of redshift and luminosity.  Our observed growth of a factor of $\\sim 2$ in satellite fraction between $z\\sim1$ and $z\\sim0$ is testament to the limited role that galaxy merging plays in galaxy evolution for $\\sim10^{12} h^{-1} M_\\odot$ mass haloes at $z<1$.  Qualitatively, our observations are consistent with a picture in which red galaxies in massive haloes have already accumulated most of their stellar mass by $z\\sim1$ and subsequently undergo little evolution until the present day.  The observed movement of the peak location for the full galaxy population is consistent with the bulk of star-formation activity migrating from higher mass haloes at high redshifts to lower mass haloes at lower redshifts.} ",
        "introduction": "\\label{sec:intro}  In our current paradigm of galaxy formation, haloes of dark matter grow from tiny imperfections in the early Universe. Against the background (at the present day) of an accelerating Universe, structures grow ``hierarchically'': small haloes form first and merge to build up larger ones, resulting in a complex filamentary network where the most massive haloes lie at the nodes of the cosmic web \\citep{2006Natur.440.1137S}.  Galaxies form as baryons fall into the centres of dark matter haloes and the inflow of cold gas provides the fuel for star formation \\citep{1978MNRAS.183..341W,1980MNRAS.193..189F}. While this scenario broadly reproduces the observed galaxy distribution over a large range of scales and cosmic time, several open questions remain concerning which processes regulate star formation inside haloes and how this leads to the Hubble sequence observed at the present day.  These physical processes which drive galaxy formation and regulate star formation are expected to correlate closely to the mass of dark matter haloes which host galaxies. Therefore, a profitable avenue to pursue to better understand the physics of galaxy formation is to determine how the relationship between the stellar mass $M_{\\rm star}$ and the dark matter halo mass $M_{\\rm h}$ changes over the lifetime of the Universe. In the local Universe, abundance-matching studies have shown that at $z<0.1$ the stellar-to-halo mass ratio (SHMR, $M_{\\rm star}/M_{\\rm h}$) reaches a maximum of a few percent at $M_{\\rm h} \\sim 10^{12}M_\\odot$, much smaller than the universal baryon fraction \\citep{2010MNRAS.404.1111G}. This indicates that the conversion efficiency of baryons to luminous galaxies is highest at these halo masses. In less massive haloes, supernovae winds are responsible for gas reheating, which reduces star formation and flattens the galaxy luminosity function at the faint end \\citep{2003ApJ...599...38B}.  In more massive haloes, feedback from active galactic nuclei \\citep{2006ApJ...652..864H,2008MNRAS.391..481S} can ``quench'' star formation, leading to formation of the ``red sequence'' of passively evolving galaxies \\citep{2006MNRAS.370..645B,2006MNRAS.365...11C}. The relative importance of these different processes as well as their dependence on environment, redshift and galaxy type is still very much an open question.  Ideally we would like to make a detailed comparison between the luminous content of dark matter haloes and the halo masses themselves as a function of redshift.  Several methodologies exist to estimate halo masses: for example, using satellite kinematics \\citep{2011MNRAS.410..210M}, X-ray temperatures \\citep{2006PhR...427....1P,2009Natur.460..213C}, or gravitational lensing \\citep{2006MNRAS.371L..60H,2006MNRAS.368..715M, 2010ApJ...724..511A,2011arXiv1104.0928L}. However, the range in redshifts over which these techniques can be applied is limited, and amassing large samples is not always easy.  At the price of making some assumptions concerning the profiles of dark matter haloes and their evolution in number density, another way of inferring halo masses is from the observed galaxy clustering and mass functions.  Usually, one makes the simplifying assumption that the number of galaxies in a given dark matter halo only depends on the halo mass, an assumption which has been verified with $N$-body simulations \\citep{2003ApJ...593....1B,2005ApJ...629..625B,2002ApJ...575..587B, 2004ApJ...609...35K}. Furthermore, the galaxy population of a halo, described by the ``halo occupation distribution'' (HOD), is assumed to be independent of environment and assembly history of the halo. Although some authors have shown that this ``halo assembly bias'' can have an observable effect on galaxy clustering measurements, the magnitude of this effect is, for the moment, much smaller than the size of systematic errors \\citep{2007MNRAS.374.1303C}. This will be an important effect to test in future surveys.  In this paper we will use this model to investigate the changing relationship of dark matter and luminous matter from $z\\sim1.2$ to the present day. Until now, the majority of papers which have used HOD modelling to interpret galaxy clustering have analysed either large, low redshift surveys such as the SDSS \\citep[see, for e.g., ][]{2010arXiv1005.2413Z,2008MNRAS.385.1257B,2010arXiv1005.2413Z}, or smaller ($\\sim1$~deg$^2$) deep fields such as the COMBO-17 and COSMOS surveys~\\citep{2006A&A...457..145P,2009MNRAS.398..807S,2011arXiv1104.0928L}. For the moment, spectroscopic surveys at higher redshifts are also limited to small fields of view~\\citep{2007ApJ...667..760Z,2010MNRAS.406.1306A}. \\cite{2008ApJ...682..937B} has interpreted clustering measurements in the framework of the halo model of slightly larger redshift baseline, although their survey only covered $\\sim7$ deg$^2$ and used five-band photometric redshifts. None of these surveys had the required combination of depth and areal coverage to make reliable measurements from low redshift to $z\\sim1$ of large samples of galaxies selected by type and intrinsic luminosity. In contrast, this is an ideal task for the Canada-France-Hawaii Legacy Survey (CFHTLS) with its unique combination of depth, area, and image quality.  The five-band $u^\\ast,g',r',i',z'$ CFHTLS photometry allows us to measure photometric redshifts out to $z\\sim1.2$ with $\\sim 3 \\%$ accuracy and small systematic errors \\citep{2009A&A...500..981C}. The large sky coverage, $\\sim133$~deg$^2$ after masking, enables us to probe a large range of halo masses. In addition, from the four independent fields of the CFHTLS we can obtain a reliable estimate of cosmic variance from the data itself. In this paper we construct a series of volume-limited samples selected by type, luminosity, and redshift, and use the halo model combined with the HOD formalism to infer how each galaxy sample populates their hosting dark matter haloes, and how this changes with look-back time.  The paper is organized as follows. In Section~\\ref{sec:observ-reduct-catal}, we present our data set and catalogue production, including data reduction and photometric redshift measurement. In Section~\\ref{sec:clustering} we outline the techniques we use to estimate galaxy clustering, and in Section~\\ref{sec:model} we present our model. Results are described in Section~\\ref{sec:results} and discussed in Section~\\ref{sec:discussion}. Conclusions are presented in Section~\\ref{sec:conclusion}.  Throughout the paper we use a flat $\\Lambda$CDM cosmology ($\\Omega_m~=~0.25$, $\\Omega_\\Lambda~=~0.75$, $H_0 = 70$~km~s$^{-1}$ ~Mpc$^{-1}$ and $\\sigma_8=0.8$) with $h~=~H_{\\rm0}/100$~km~s$^{-1}$~Mpc$^{-1}$. All magnitudes are given in the AB system. ",
        "conclusions": "\\label{sec:conclusion}  We have made one of the most precise measurements to date of the angular correlation function $w$ and its dependence on luminosity and rest-frame colour from $z\\sim0$ to $z\\sim1$. Our measurements were constructed from a series of $i'<22.5$ volume-limited samples, containing $\\sim 3\\times 10^6$ galaxies in four independent fields. These samples were created using accurate five-band photometric redshifts in the CFHTLS Wide survey calibrated with $\\sim10^4$ spectroscopic redshifts. A cross-correlation analysis showed consistency with a relatively small contamination fraction arising from scatter between photometric redshift bins.  We interpreted these measurements in the framework of a model in which the number of galaxies inside each dark matter halo is parametrised by a simple analytic function (the halo occupation distribution) with five adjustable parameters. In this model, central and satellite galaxies are treated separately. For each galaxy sample, we performed a full likelihood analysis and explored model parameter space using an efficient Population Monte Carlo analysis. The very large survey volume probed by the CFHTLS Wide allowed us to place robust constraints on the redshift evolution of the halo model parameters, and make accurate estimates for halo properties of many different galaxy samples over more than two orders of magnitude in luminosity and three orders of magnitude in halo mass, from $10^{12}$ to $10^{15}(h^{-1} M_{\\sun})$.  Using 30-band photometric data from the COSMOS survey, we derived an empirical relation between stellar mass, luminosity and redshift and used this to convert our luminosity-selected samples to stellar-mass limited samples. All of our conclusions (following below) are independent of this correction.  These are our principal results: \\begin{enumerate} \\item For a given luminosity threshold, galaxies with redder rest-frame colours are more clustered than bluer ones. Clustering strength increases with increasing luminosity, reflecting that bright galaxies reside in more strongly clustered massive haloes.  \\item We consider the redshift and luminosity evolution of halo parameters $M_1$, representing the characteristic halo mass required to host at least one satellite galaxy, and $M_{\\rm min}$, representing the halo mass where on average 50\\% of haloes contain one galaxy. For our full sample, $M_1$ is closely approximated by $\\sim 17\\times M_{\\rm min}$. This ratio becomes smaller at higher redshifts for lower-mass haloes. For red galaxies, $M_1/M_{\\rm min}$ is $\\sim12$ at intermediate mass scale, and remains constant for all halo masses over $0.2 < z < 0.8$. This lower ratio for redder galaxies reflects the higher abundance of satellites in these haloes; in general, redder populations reside in more massive haloes.  \\item By fitting a simple analytic relation between central galaxy luminosity and halo mass to our observations, we find a maximum in the stellar-to-halo mass ratio. In the lowest redshift bin $z\\sim0.3$, this peak mass is $M^{\\rm peak}_{\\rm h} = 4.5\\times 10^{11} h^{-1} M_{\\sun}$ for the full samples, and $M^{\\rm peak}_{\\rm h} =21 \\times 10^{11} h^{-1} M_{\\sun}$ for the red samples, respectively. This transition represents the halo masses where baryons were most efficiently converted into stars, and is in good agreement with measurements from other studies.  \\item For the full sample, $M^{\\rm peak}_{\\rm h}$ shifts to higher halo masses at higher redshifts. For the red galaxy sample, the peak position evolves less rapidly with redshift. These results can be understood qualitatively from the lack of on-going star-formation in the red galaxy population which means that the stellar mass content in these massive haloes changes very slowly. For the full galaxy sample, which is expected to contain galaxies still undergoing active star-formation, the shift of the transition mass to progressively more massive haloes at higher redshifts is a manifestation of ``anti-hierarchical'' galaxy formation. These results also indicate that the massive, passive galaxies in our survey are already fully formed by $z\\sim1.2$.  \\item At increasing luminosities, the fraction of satellite galaxies rapidly drops to zero. For less luminous galaxies, the satellite fraction is higher in our red sample than in our full sample, reflecting the high number of faint red satellites, consistent with observations in the local Universe \\citep{2006MNRAS.368..715M,2010arXiv1005.2413Z}. For our full sample, the number of satellites increases by a factor of two from $z\\sim1$ to $z\\sim0$, consistent with a combined study of galaxies in DEEP2 and SDSS \\citep{2007ApJ...667..760Z}. This result suggests that the occupation function for satellite galaxies inside intermediate-mass haloes ($\\sim10^{12} h^{-1} M_{\\sun}$) remains constant with cosmic time, and implies that the galaxy merging rate does not change since $z=1$.  \\item Above the characteristic sample luminosity, the mean galaxy bias $b_{\\rm g}$ increases rapidly. Faint red samples have $b_{\\rm g}\\sim1.5$, compared to $b_{\\rm g}\\sim1$ for the full sample. Redder samples also have higher average halo masses, which decrease towards higher redshifts.  \\item We compare the measured evolution of our mean halo masses, biases and satellite fractions with a model assuming constant HOD parameters over our redshift range.  This model provides a remarkably good fit to our measurements for massive haloes. We note that lower mean halo masses (particularly in the full sample) show a slower decrease compared to a redshift-independent HOD model.  \\end{enumerate}  Systematic photometric redshift errors have little impact on our faint luminosity-threshold samples, for which we find no evidence for incompleteness larger than the field-to-field variance estimates.  Our very brightest samples however may be affected by incompleteness or contamination from fainter objects. We have shown that these errors could reduce our measured clustering amplitudes and consequently could lead to an underestimation of halo masses. The effect will be more important in contaminated samples, where the higher galaxy number density will lead to lower halo masses. In the worst case, $\\log M_{\\rm min}$ might be underestimated by $\\sim0.5$.  Because of the lack of deep near-infrared data in the CFHTLS, we cannot calculate reliable stellar masses and cleanly separate the active and passive populations in mass-selected samples. Such a data set, perhaps combined with approximate local density estimators calibrated using spectroscopic redshifts, and a refined halo model would allow us to understand in greater detail the physical origin of the shift of the transition masses to higher halo masses observed here.  Furthermore, given the fact that the peak of the star-formation and mass assembly takes place in the redshift range $1<z<2$, it is important to push these studies to higher redshifts by constructing larger, wide-area surveys with deep near-infrared data.  In our model, we have assumed a halo mass function, halo density profile, and a cosmological model, although it is well known that the evolution of galaxy clustering depends on these parameters. Combining galaxy-galaxy lensing, galaxy clustering and mass-function measurements from an expanded CFHTLS including near-infrared data should provide new insights on both cosmology and galaxy evolution."
    },
    "1107/1107.3152_arXiv.txt": {
        "abstract": "We have analyzed the data on 16,836 RR Lyrae (RR Lyr) variables observed toward the Galactic bulge during the third phase of the Optical Gravitational Lensing Experiment (OGLE-III), which took place in 2001--2009. Using these standard candles, we show that the ratio of total to selective extinction toward the bulge is given by $R_I=A_I/E(V-I)=1.080\\pm0.007$ and is independent of color. We demonstrate that the bulge RR Lyr stars form a metal-uniform population, slightly elongated in its inner part. The photometrically derived metallicity distribution is sharply peaked at $\\feh=-1.02\\pm0.18$, with a dispersion of 0.25~dex. In the inner regions ($|l|<3\\degr$, $|b|<4\\degr$) the RR Lyr tend to follow the barred distribution of the bulge red clump giants. The distance to the Milky Way center inferred from the bulge RR Lyr is $R_0=8.54\\pm0.42$~kpc. We report a break in the mean density distribution at a distance of $\\sim0.5$~kpc from the center indicating its likely flattening. Using the OGLE-III data, we assess that $(4$--$7) \\times 10^4$ type ab RR Lyr variables should be detected toward the bulge area of the on going near-IR VISTA Variables in the Via Lactea (VVV) survey, where the uncertainty partially results from the unknown RR Lyr spatial density distribution within $0.2$~kpc from the Galactic center. ",
        "introduction": "\\label{sec:introduction}  RR Lyrae (RR Lyr) variable stars are known as standard candles and useful tracers of old populations. They play an essential role in our understanding of the formation and evolution of the Galaxy. When multi-epoch light curves are available, RR Lyr variables can be easily identified. A large number of observations allows determination of their properties with high precision. The derived parameters can be used to determine interstellar extinction, metallicity, distance, and spatial distribution of the stars. RR Lyr stars observed toward the bulge provide an independent determination of the distance to the center of the Milky Way.  RR Lyr stars observed close to the central regions of the Milky Way concentrate toward the Galactic center was first noticed by \\cite{1932BAN.....6..163V,1933BAN.....7...21V}. Later, \\cite{1946PASP...58..249B} found a strong predominance of RR Lyr variables toward a relatively unobscured area today called the Baade's window, indicating the presence of Population II stars in the Milky Way center. By the early 1990s approximately 1000 RR Lyr variables inhomogeneously distributed toward the bulge were known.  The number of new bulge RR Lyr variables has increased following the advent of massive photometric surveys. \\cite{1994AcA....44..317U, 1995AcA....45....1U, 1995AcA....45..433U, 1996AcA....46...51U, 1997AcA....47....1U} published a catalog of over 3000 periodic variable stars detected in the fields covered by the first phase of the Optical Gravitational Lensing Experiment (OGLE-I). Two-hundred and fifteen of these objects were classified as RR Lyr. Analysis of the OGLE-II data brought a much larger list of 2713 RR Lyr (\\citealt{2003AcA....53..307M}). Later, using the same source of data, \\cite{2006ApJ...651..197C} prepared a catalog of 1888 fundamental mode RR Lyr stars (type RRab).  The MACHO microlensing project also observed a numerous sample of RR Lyr stars toward the Galactic center. \\cite{1998IAUS..184..123M}, using a sample of 1150 RRab and 550 RRc stars, showed that the spatial distribution of RR Lyr between 0.3 and 3~kpc follows a power law with an inclination of $-3.0$. \\cite{1998ApJ...492..190A} examined the mean colors and magnitudes of $\\sim$1800 RR Lyr and found that the bulk of the population is not barred. Only RR Lyr located toward the inner fields closer to the Galactic center ($l<4\\degr$, $b>-4\\degr$) seem to follow the barred distribution observed for red clump giants (RCGs) (\\citealt{1994ApJ...429L..73}). Recently, \\cite{2008AJ....135..631K} analyzed photometric data on 3525 MACHO RRab stars to assess the reddening toward the Galactic bulge. They derived the selective extinction coefficient $R_{V,VR}=A_V/E(V-R)=4.3\\pm0.2$, which corresponds to the average value observed in the solar neighborhood $R_{V,BV}=A_V/E(B-V)=3.1\\pm0.3$  (\\citealt{1989ApJ...345..245C,1999PASP..111...63F}).  The first metallicity measurements of bulge RR Lyr stars were made by \\cite{1976ApJ...210..120B} using the $\\Delta S$ method (\\citealt{1959ApJ...130..507P}). Using 9 stars in the ~100 variable sample of \\cite{1963Sci...140..658I} list, they obtained $\\langle\\feh\\rangle$ $=-0.65\\pm0.15$~dex and concluded that the stars are mildly metal-poor. \\cite{1986A&A...169..111G} used a sample of 17 bulge RR Lyr variables and found a wide range of iron abundances, between $-1.8$ and $+0.1$~dex. Later, from spectra of 59 RRab and RRc variables, \\cite{1991ApJ...378..119W} determined an average metallicity of $\\langle\\feh\\rangle$ $=-1.0$~dex on the \\cite{1984ApJS...55...45Z} metallicity scale. In contrast to previous studies they found the metallicity distribution to be sharply peaked, with a dispersion of only 0.16~dex. The most recent metallicity estimate was performed by \\cite{2008AJ....136.2441K}. Based on photometric data for 2,435 MACHO RRab stars they derived $\\langle\\feh\\rangle$ $=-1.25$ (on the \\citealt{1984ApJS...55...45Z} scale), with a broad metallicity range from $\\feh=-2.26$ to $-0.15$~dex.  The bulge RR Lyr stars have also often been used to measure the distance to the Galactic center, under the assumption that the center of their population corresponds to the center of the Milky Way. The first such measurement by \\cite{1951POMic..10....7B} yielded $R_0=8.7$~kpc. For more than half a century, various investigations brought different values of $R_0$ with comparable uncertainty, e.g., \\cite{1972ApJ...174..573H} obtained $R_0=7.0\\pm0.6$~kpc, \\cite{1975A&A....41...71O} $R_0=8.7\\pm0.6$~kpc, \\cite{1985MmSAI..56...15B} $R_0=6.94\\pm0.58$~kpc, \\cite{1986MNRAS.220...69W} $R_0=8.1\\pm0.4$~kpc, \\cite{1987MNRAS.226..927F} $R_0=8.0\\pm0.65$~kpc, \\cite{1995AJ....110.1674C} $R_0=7.8\\pm0.4$~kpc, \\cite{1997MNRAS.284..761F} $R_0=8.1\\pm0.4$~kpc, \\cite{2006ApJ...651..197C} $R_0=8.3\\pm0.7$~kpc, \\cite{2008A&A...481..441G} $R_0=7.94\\pm0.63$~kpc, and recently \\cite{2010AcA....60...55M} $R_0=8.1\\pm0.6$~kpc.  In this paper, using the newly released catalog (\\citealt{2011AcA....61....1S}) of nearly 17,000 RR Lyr stars detected toward the Galactic bulge during the third phase of the OGLE project (OGLE-III, \\citealt{2003AcA....53..291U,2008AcA....58...69U}), we show that these variables constitute a uniform population, different to other populations residing in the central regions of our Galaxy.  The structure of this paper is as follows. Section~\\ref{sec:data} describes the selection of the sample and gives information on the reddening. Metallicities and magnitudes of the bulge RR Lyr variables are determined in Section~\\ref{sec:metalmags}. In Section~\\ref{sec:distance}, we estimate the distance to the Galactic center. The structure of the RR Lyr population, based on their mean brightness and density distribution, is studied in Section~\\ref{sec:structure}. Section~\\ref{sec:conclusions} states our main conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have analyzed the data on 16,836 RR Lyr stars observed toward the Galactic bulge during the third phase of the OGLE project (years 2001--2009). After eliminating low-amplitude stars, objects likely foreground and background sources, and GC members the sample includes 10,472 RRab, 4608 RRc, and 78 RRd bulge variables with derived mean $V$- and $I$-band brightness. For an additional 627 RRab, 151 RRc, and 2 RRd stars, all located close to the Galactic plane, we possess mean magnitudes in the $I$ filter alone.  Using color information for this type of standard candles, we find that the ratio of total to selective extinction $R_I=A_I/E(V-I)$ is independent of color and equals $1.080\\pm0.007$. This result confirms the anomalous nature of the extinction toward the bulge.  Our analysis clearly demonstrates that the bulge RR Lyr stars constitute a very uniform population, different to the majority of the stars in the central regions of the Milky Way. The photometrically derived metallicity distribution is sharply peaked at $\\feh_{\\rm J95}=-1.02\\pm0.18$ with a dispersion of 0.25~dex on \\cite{1995AcA....45..653J} metallicity scale. Based on the distribution of dereddened brightness we show that the RR Lyr population is elongated in the inner part, where it tends to follow the barred distribution of the bulge RCGs. The elongated structure can also be noticed as a flat peak in the distance distribution. A possible explanation of this structure along the bar is that less populated old RR Lyr stars may feel gravitational forces from young stars in the Milky Way bulge. Farther off the Galactic plane (for $|b|\\gtrsim4\\degr$) the shape of RR Lyr population becomes spherical. We note that, due to small number statistics, we can not conclude the presence of a split in the population of bulge RR Lyr stars, as is observed in RCGs.  The distance to the Milky Way center inferred from the OGLE-III bulge RRab variables is $R_0=8.54\\pm0.42$~kpc. This result is in agreement with $R_0=8.1\\pm0.6$~kpc obtained from the OGLE-II data (\\citealt{2010AcA....60...55M}) and with a value of $R_0=8.15\\pm0.49$~kpc combined by \\cite{2010RvMP...82.3121G} from 11 different methods. It seems that the RR Lyr method suffers from uncertainties in the metallicity determination, and hence absolute brightness of the stars.  In our work, we have also studied the density distribution of the bulge RR Lyr stars as a function of distance from the Galactic center. We found a break in the distribution at a distance of $\\sim0.5$~kpc indicating flattening toward the center. Taking into account completeness of the OGLE-III data, we estimate the number of RRab stars which should be detected within the bulge area of the near-IR VVV survey as $(4$--$7)\\times10^4$ stars. The uncertainty partially results from unknown density distribution within $0.2$~kpc from the center.  On going and future wide-field photometric surveys, such as OGLE-IV, VVV, Pan-STARRS (\\citealt{2002SPIE.4836..154K}), and Large Synoptic Survey Telescope (\\citealt{2008ApJ...684..287I}) should bring definitive answers to questions pertaining to the structure and evolution of all Galactic RR Lyr populations. For example, one of unsolved issues is whether the bulge and halo RR Lyr variables form the same population. \\cite{1998IAUS..184..123M} suggested that the bulge RR Lyr stars could represent the inner extension of the Galactic halo. The bulge pulsators do not show the \\cite{1939Obs....62..104O} dichotomy observed in the halo variables. From the period distribution and period--amplitude diagram presented in \\cite{2011AcA....61....1S} it seems that probably all bulge RR Lyr stars belong to the Oosterhoff type I (OoI). Our work shows that the bulge population is very metal uniform and was likely formed on a short time scale. The Galactic halo was probably formed by at least two distinct accretion processes (\\citealt{2008ApJ...678..865M}). About 25\\% of the halo RR Lyr variables observed at distances 3--30~kpc appear to be Oosterhoff type II (OoII) stars with a significantly different radial density profile to the remaining 75\\% of OoI objects. The OoI halo population has a power-law exponent of $-2.26\\pm0.1$ while the OoII component has a steeper slope with $-2.88\\pm0.04$ (\\citealt{2008ApJ...678..865M}). The OoI stars are on average slightly more metal-rich ($\\langle\\feh_{\\rm J95}\\rangle\\sim-1.7$~dex) than the OoII stars ($\\langle\\feh_{\\rm J95}\\rangle\\sim-2.0$~dex; \\citealt{2010ApJ...708..717S}). Comparison of the above numbers for the halo RR Lyr variables with those for the bulge objects obtained in this work does not allow us to answer the question on the common bulge/halo population. More data, in particular on RR Lyr stars residing in the central and outer regions of our Galaxy, are needed."
    },
    "1107/1107.1966_arXiv.txt": {
        "abstract": "Galaxies with polar rings (PRGs) are a unique class of extragalactic objects allowing to investigate a wide range of problems, linked with the formation and evolution of galaxies, and to study the properties of their dark haloes. The progress in the study of PRGs is constrained by a small number of known objects of this type. Whitmore et al. (1990) compiled a catalogue and photographic atlas of PRGs and related objects that included 157 galaxies. Up to date, we can only attribute about two dozens of kinematically-confirmed galaxies to this class, mostly from this catalogue.  We present a new catalogue of PRGs, supplementing the catalogue of Whitmore et al. (1990), and significantly increasing the number of known candidate PRGs. The catalogue is based on the results of the original Galaxy Zoo project, under which the volunteers performed a visual classification of nearly a million galaxies from the SDSS. Based on the preliminary classification of the Galaxy Zoo, we viewed more than 40\\,000 images of the Sloan Digital Sky Survey (SDSS) and selected  275 galaxies, included in our catalogue.  Our Sloan-based Polar Ring Catalog (SPRC) contains  70 galaxies that we classified as ``the best candidates'', among which we expect to have a very high proportion of true PRGs, and  115 good  PRG candidates.   53 galaxies are classified as PRG related  objects (mostly, the galaxies with strongly warped discs, and mergers). In addition, we identified  37 galaxies that have their presumed polar rings strongly inclined to the line of sight (seen almost face-on).  The SPRC objects are on the average fainter and located further away than the galaxies from the catalog by Whitmore et al. (1990), although our catalogue does include dozens of new nearby candidate PRGs.  The new catalogue significantly  increases the number of genuine PRG candidates, and may serve as a good basis both for the further detailed study of individual galaxies, and for the statistical analysis of PRGs as a separate class of objects.  We performed spectroscopic observations of six galaxies from the SPRC at the 6-m BTA telescope. The existence of polar rings was confirmed in five galaxies, and one object appeared to be a projection of a pair of galaxies. Adding the literature data, we can already classify 10 galaxies from our catalogue to the kinematically-confirmed PRGs. ",
        "introduction": "Galaxies with polar rings (PRGs) are an interesting case of peculiar systems that reveal outer rings or discs of gas, dust and stars, rotating in the plane approximately perpendicular to the disc of the main (or host) galaxy. It is believed that the formation of PRGs is in most cases caused by galaxy mergers with the corresponding direction of angular momentum, the accretion by the host galaxy of the companion's matter, or gas filaments from the intergalactic medium  \\citep[see references in the review by ][]{Combes06}. It is shown that in the case of an oblate or triaxial gravitational potential, stable orbits exist in the polar plane, hence, the captured material of the companion will be rotating here long enough. If the orbital plane is notably different from polar, then the formed ring would relatively quickly precess to the Galactic plane \\citep{Steiman1982}.  A detailed examination of the structure and dynamics of PRGs both clarifies the role of interactions and mergers in the current evolution of galaxies, and also allows to study the features of the distribution of gravitational potential at large radii. Since we are observing here a circular rotation in two mutually perpendicular planes, it becomes possible to study the three-dimensional distribution of mass in the galaxy. Particularly we can determine the shape of the dark halo: its oblateness, elongation, and deviations from axial symmetry \\citep{SackettPogge1995,Iodice03,Combes06}.  In addition, there are recent indications that some relatively massive polar rings are formed by a cold accretion of gas from the filaments of the intergalactic medium. In this scenario, \\citet*{Maccio2006} and \\citet{Brook2008} were able to construct an object very similar to the well-known PRG NGC~4650A using cosmological hydrodynamic simulations. Recently, additional observational arguments supporting the model of accretion from gas filaments were presented.  For instance, \\citet{Spavone2010} suggested that NGC~4650A has the gas metallicity lower than a spiral galaxy disc of the same total luminosity. Also, \\citet{Stanonik2009} have found in the galaxy SDSS J102819+623502 a polar disc of neutral hydrogen aligned nearly perpendicularly to a cosmological wall situated between two voids. Therefore, a further study of PRGs will help to better perceive the causes (and the source) of the accretion of gas from the intergalactic medium, and will help to resolve some cosmological puzzles of galactic evolution \\citep*{Combes2008}.%   Although the first credible evidences of the existence of polar rings were obtained quite long ago \\citep{SchechterGunn1978}, their statistical study began only after \\citet{Whitmore1990} have published a catalogue of PRG candidates -- the Polar Rings Catalog=PRC. In total, their catalogue contained 157 objects, from which only six (PRC A category) already had a kinematic confirmation (rotation, detected in two orthogonal planes), and 27 galaxies were included the category of `good candidates' (PRC B). The greater part of the list was occupied by 73 ``possible candidates'' (PRC C) and 51 ``related objects'' (PRC D). Over the following twenty years, several groups were involved in the study of kinematics of catalogue objects (mostly -- in the northern celestial hemisphere) both using the methods of optical spectroscopy and radio interferometry. However, the number of kinematically confirmed PRGs is still meager. In particular, the outer polar rings were confirmed in about 20 galaxies from the PRC \\citep*[]{Resh2011}, including several ``related systems'', where the polar structures are only forming. In addition, six other galaxies contain inner polar discs with radii smaller than $1-3$ kpc.    There are even fewer PRGs studied in detail, we can currently speak here of less than a dozen galaxies only. The question of the shape of their dark halo remains open. On the one hand, the statistical analysis of the distribution of integral parameters of PRGs in the Tully-Fisher diagram indicates that the outer halo is flattened to the plane of the polar ring \\citep{Iodice03,Resh2004}. This means that the major principal plane of the halo is orthogonal to the disc of the central galaxy, which in itself is unusual. On the other hand, the results of a study of several galaxies with strongly warped HI discs, such as NGC~2685 = PRC A-3 \\citep{Jozsa2009}, NGC~3718=PRC D-18 \\citep{Sparke2009}, and NGC~4753=PRC D-23\\citep*{Steiman1992} indicate a more or less spherical halo.  The complexity of the problem and the need for a detailed analysis of each galaxy is illustrated by the case of NGC4650A. \\citet*{Whitmore1987} have earlier noted that the gravitational potential of the galaxy is almost spherical, but the analysis of  kinematics of more extended structures indicates a very strong oblateness of the halo, oriented in accordance with the polar ring \\citep{Iodice2010}.    The gravitational stability of polar rings is a subject of debate as well. \\citet{Wakamatsu1993} shows that while the polar ring gas clouds pass through the gravitational well of the stellar disc, the gas generates shock waves that ``clean out'' the inner region of the ring. Obviously, they can additionally induce star formation in the ring. Numerical hydrodynamic models by \\citet*{Theis2006} also indicate the possibility of formation of a stable spiral structure  in the polar ring of   NGC4650A. Unfortunately, the observational evidence for the existence of spirals in the polar rings is yet rare and contradictory, since to confirm this hypothesis, the polar disc has to be sufficiently inclined to the line of sight. In addition, there are no reliable observational estimates of the Toomre's parameter $Q$ for the polar rings and discs as yet.   There are PRGs, discovered outside the catalogue by \\citet{Whitmore1990}.  This is a dwarf galaxy NGC~6822 of the Local Group, in which the outer HI disc rotates in the plane, polar to the stellar disc \\citep{Demers2006}.   A similar but larger structure is detected around a more distant galaxy SDSS J102819.24+623502.6 \\citep{Stanonik2009}. Ultraviolet images of  NGC~4262 reveal signatures of star formation in the gas ring surrounding it  \\citep{Bettoni2010}. The most  distant confirmed published  polar ring ($z=0.06$) was studied by \\citep{Brosch2010}.  In the Hubble deep fields three even more distant PRG candidates were discovered at  $z=0.6-1.3$, which are still awaiting confirmation \\citep{Resh1997,Resh2007}.   Two directions can be envisioned in the prospects for the  further studies of PRGs. First of all, a detailed study of the known candidates and their environments needs to be done using the data on the morphology and kinematics in different spectral ranges, from the UV to radio. Secondly, we have to expand the list of candidates, to both clarify the luminosity function of PRGs, and to move towards higher redshifts, but as well to search for objects in which both the ring and the galaxy are ``conveniently'' oriented to the line of sight, allowing to simultaneously study their kinematics and structural details. The PRC was based on the study of photographs of individual galaxies. In the modern era, it is reasonable to use digital sky surveys, such as the SDSS for these purposes.   Note that we do not discuss in this paper the so-called inner polar rings and discs, observed in the circumnuclear regions of nearby, typically early-type galaxies \\citep[see references and discussion in] [] {Corsini2003, SilAfan2004}. Such structures \\citep[$\\sim30$ of which are already known, see][]{Moiseev2010} are lost in the bright background of the bulge, and are discovered mainly due to their kinematics, decoupled from the  galaxy disc.  In this paper, we present a new list of PRG candidates, a few of which have already been confirmed. Section  \\ref{sec_catalog} describes the technique of catalogue compilation using the data from the  Galaxy Zoo project. Section \\ref{sec_ABCD} describes the division of catalogue objects into several types.  In Section \\ref{sect_obs} we  present the information about five galaxies for which  there already exist detailed studies of the internal kinematics. We managed as well to perform spectral observations of six galaxies with the 6-m BTA telescope of the SAO RAS. Five of the observed objects were confirmed to be classical PRGs, and one turned out to be a projection of an interacting pair of galaxies. Section \\ref{sec_conclusion} briefly discusses the results of this paper.    \\begin{figure} \\includegraphics[width=0.45\\textwidth]{EL_c-eps-converted-to.pdf} \\includegraphics[width=0.45\\textwidth]{CS_c-eps-converted-to.pdf} \\includegraphics[width=0.45\\textwidth]{EDGE_c-eps-converted-to.pdf} \\includegraphics[width=0.45\\textwidth]{Merging_c-eps-converted-to.pdf} \\includegraphics[width=0.45\\textwidth]{unknown_c-eps-converted-to.pdf} \\caption{Histograms of distribution of relative galaxy numbers within each type. The blue line -- for all the  893,212 galaxies of the Galaxy Zoo, red -- for 126 PRG candidates from the reference sample. Dashed lines show the selection criteria.} \\label{fig_zoo} \\end{figure} ",
        "conclusions": "\\label{sec_conclusion}   As a result of first spectral observations of galaxies from our catalogue we were able to confirm the presence of polar rings in five of them. The sixth object, SPRC-178, turned out to be a pair of galaxies. All the confirmed candidates reveal both ionized gas and stars in their polar orbits. We hope that a more in-depth analysis of these and subsequent observations will allow us to say more about the stellar population of   polar rings: their age, metallicity and kinematics. Note that by now we  only knew of three galaxies, in which the motion of stars in polar rings was studied -- NGC~4650A \\citep{SwatersRubin2003}, UGC~5119 \\citep{Merkulova2008}, and UGC~2748 \\citep{Merkulova2009}.   Along with the available literature data, the kinematic confirmation of the presence of an outer polar component exists for 10 galaxies from our catalogue. This is already a significant supplement to the number of known PRGs.    Supplementing the PRC galaxies with the objects from our SPRC catalogue, we have by 3 times increased the number of known objects that are in one way or another connected with the phenomenon of polar rings. Thereby, the number of ``genuine'' PRG candidates has increased almost threefold, since the 33 PRC-A and PRC-B objects were subjoined by   70 best candidates  from the SPRC. Fig.~\\ref{fig_z} compares the redshift distributions of galaxies from \\citet{Whitmore1990} and SPRC. It is evident that the new catalogue contains a greater number of distant objects, starting from $z>0.05$. A sharp decline in the number of candidates at $z>0.1$ is obviously connected with the selection -- the images obtained with $seeing\\approx1$ arcsec allow to distinguish structural features only in   galaxies having relatively large angular sizes. At the same time, dozens of new nearby ($z<0.03$) PRG candidates were discovered.    Therefore, even among the nearby galaxies, the PRGs occur several times more frequently than it was previously thought. By a rough estimate of \\citet{Whitmore1990}, only 0.5 per cent of nearby S0 galaxies possess polar rings, while this estimate could be significantly increased in view of the rings, seen face-on. More accurate recent estimates by \\citet{Resh2011}, based on the constructed luminosity function, prove that the rate of PRG occurrence among the nearby galaxies   in the range of $M_B=-17^m...-22^m$ amounts to 0.13 per cent, and adjusted for the effect of projection it reaches  0.4 per cent. Employment of new catalogue objects can significantly increase this value, since, according to the criteria, described in Section \\ref{sect_zoo}, the actual number of PRGs among all the Galaxy Zoo galaxies is about 3 times greater than in our catalogue. Moreover, the PRC was based on viewing the photographic plates with a smaller limit on surface brightness, than the SDSS data have. This is why the objects with fainter outer structures made it into the SPRC catalogue. However, before reviewing the luminosity function, we have to make the kinematic confirmation of new candidates.      The galaxies with massive outer polar rings are of particular interest. Formation of these rings can not always be explained within the classical scenario of accretion or merging. The best  example known to date is a nearby galaxy NGC~4650A, possessing a massive self-gravitating stellar-gaseous ring. A recent paper by \\citet{Spavone2010} hypothesizes that such a polar disc in this galaxy was formed by accretion of gas from the intergalactic medium filaments. Among the objects of our catalogue there are several galaxies with bright and extended polar rings: SPRC-7, SPRC-27, and SPRC-40. \\citet{Brosch2010} have shown that the former galaxy can largely be regarded as a giant analogue of NGC~4650A. Further detailed study of similar systems will help to better understand the processes of gas accretion  during the galaxy formation on the cosmological scales.     \\begin{figure} \\includegraphics[width=0.45\\textwidth]{fig_z_col-eps-converted-to.pdf} \\caption{Histograms of the distribution of PRG candidates by redshift for the \\citet{Whitmore1990} catalogue and the new list.} \\label{fig_z} \\end{figure}"
    },
    "1107/1107.1157_arXiv.txt": {
        "abstract": "{ We present the design and performance of a 2x2 prototype array of corrugated feed-horns in W-band. The module is fabricated using a so-called ``platelet'' technique by milling Aluminum plates. This technique is suitable for low-cost and scalable high performance applications. Room temperature Return Loss measurements show a low (<-30 dB) reflection over a 30\\% bandwidth with a maximum matching of -42 dB at 100 GHz for all four antennas. Beam pattern measurements indicate good repeatability and a low (-25 dB) sidelobe and crosspolarisation levels. This work is particularly relevant for future Cosmic Microwave Background polarisation measurements, which require large microwave cryogenic detector arrays coupled to high performance corrugated feed horns. } ",
        "introduction": "Forty-five years after its discovery, the \\textbf{C}osmic \\textbf{M}icrowave \\textbf{B}ackground (CMB) radiation is today a very powerful tool of cosmological investigation.  The CMB originates from the decoupling of matter and radiation, occurred when the temperature of the expanding universe fell below 3000\\,K. At this temperature the plasma could combine into neutral hydrogen, thus reducing the Thomson scattering cross section and letting the radiation propagate freely. Because of cosmic expansion the CMB can be observed today as a blackbody radiation at the temperature T=2.725 $\\pm$ 0.002 K \\cite{Mather} peaked in the microwave region of the electromagnetic spectrum.  The CMB temperature is anisotropic at the level of $\\Delta T/T\\approx10^{-5}$ on scales from 180 degrees to about $5^{\\prime}$. The angular temperature anisotropies depend on the initial conditions of the Universe and on the physical processes that occurred from the Big Bang to the matter-radiation decoupling and later via Sunyaev-Zel'dovich effect. Accurate measurements of these anisotropies, therefore, provide a wealth of information on the universe history and composition (baryons, dark matter, dark energy). CMB temperature anisotropies have been measured by many ground and balloon experiments as well as by the NASA satellites COBE\\footnote{\\url{http://aether.lbl.gov/www/projects/cobe/}} and WMAP\\footnote{\\url{http://map.gsfc.nasa.gov/}} (see  \\cite{Anisotropies} for an extensive review until WMAP). The third generation ESA satellite Planck\\footnote{\\url{http://planck.esa.int}}, launched in May 2009, will provide in few years foreground-limited CMB temperature anisotropy measurements and open the path to accurate measurements of CMB polarisation anisotropies, the future of CMB experiments.  Polarisation anisotropies, generated by Thomson scattering, can be represented on the celestial sphere by two modes, named E-mode and B-mode, according to their symmetry properties. The detection of B-modes would provide an indirect signature of cosmological gravitational waves, that are predicted by the Hot Big Bang model \\cite{Pol_Primer}.  Many experiments are currently planned to accurately measure the CMB polarisation. First detection of E-modes and its correlation with the temperature anisotropy was published in 2003 by the DASI collaboration \\cite{DASI_and_others}, while the most up-to-date measurements come from WMAP \\cite{WMAP_pol}. Because only 10\\% of the CMB is polarised, polarisation anisotropies are smaller by a factor 10 compared to temperature anisotropies. E-modes, in particular, are at the $\\mu$K level, while for B-modes the signal is expected to be much weaker and not even well theoretically constrained, although it is not exptected to be larger than some fraction of $\\mu$K \\cite{Pol_Primer}.  Detection of such small signals call for ultra-high sensitivity instruments constituted of large detector arrays of hundreds and even thousands of channels. A key element is represented by the availability of high performance antennas which represent the interface between the sky (or telescope) and the receiver. Precision polarisation measurements require optical components with highly symmetrical optical response and losses not greater than $\\sim$1\\%. Corrugated feed-horns match these requirements and are commonly used in most CMB experiments [cfr. WMAP and Planck]. Typical manufacturing techniques include direct machining, up $\\sim$40 GHz, and electroforming, at higher frequencies. These methods can be limited by time and cost when applied to the production of large detector arrays of hundreds of feed-horns, as requested by future CMB polarisation experiments, so cheaper and quicker techniques are required.  In this work we investigate the validity of a technique, named ``platelet'' \\cite{Kangas}, \\cite{Haas}, \\cite{Wollack}, for the mass production of corrugated feed-horn arrays. In particular we focus on three key factors: (i) mechanical feasibility with reduced costs and time, (ii) reliability in standard and cryogenic conditions and (iii) repeatability of the electromagnetic horn performance. Recently the QUIET\\footnote{\\url{http://quiet.uchicago.edu}} experiment started CMB polarisation measurememnts using, for the first time, feed-horn array realized with the platelet technique \\cite{QUIET2010}.  The platelet technique consists in the overlap of properly machined metal sheets so that the feed-horn is then realized in layers. This method is relatively cheap, with a cost saving of about 90\\% with respect to the electroforming and allows quick manufacturing of feed-horn arrays by drilling several holes in each plate. A complete corrugation, with a tooth and a groove, can be obtained in each plate. A complex focal surface can be obtained with a series of modules properly tilted to reproduce the required geometry. Particular care is required in the plate assembly alignment system and for mass/volume control in order to minimize the weight.  The advantages of this technique go beyond its application to CMB measurements and can be exploited in a wide range of mm-wave and sub-mm astronomical instruments. In particular, platelet techniques may prove appropriate for the new generation of large interferometric arrays, such as ALMA, which require high performance antennas (including polarisation properties) over a large number of channels. ",
        "conclusions": "In this paper we have discussed the manufacturing and testing of a non-bonded W-band 2$\\times$2 array prototype of corrugated feed-horns realised by ``platelet'' technique. Feed-horns have been realized by layers, allowing reduced times and costs with respect to standard techniques like electroformation, typically adopted for frequencies $\\geq$ 100~GHz.  Beam patterns and RL measurements are in good agreement with the corresponding simulations and show high repeatability for all the feed-horns, also in cryogenic conditions. In particular the RL is less than -30\\,dB over a 30\\% bandwidth with a maximum matching of -42\\,dB at 100\\,GHz for all four antennas. Also beam pattern measurements indicate very good repeatability and a low level of side lobes and cross-polar polarisation (both lower than -25\\,dB). Moreover an estimation of the IL indicates that it is not worse than -0.06 dB over the W band.  These results are very promising for the time and cost-effective mass production of large array of high-performances corrugated feed-horns, as showed by \\cite{Wollack}. The non-bonded solution, chosen here and in \\cite{Kangas} to assemble the module, allow simple manufacturing with the possibility to disassemble the module if needed and maintain electromagnetic performance at cryogenic conditions. Further developments of this assembling solution are needed (and in progress) for increasing the packing efficiency (horn aperture area/total area) to 0.3 and reducing the mass. These developments will enable the realization of large, lightweight and high density focal planes, as required by the next generation experiments for the precise measurements of the CMB polarisation."
    },
    "1107/1107.4772_arXiv.txt": {
        "abstract": "The prospect of detecting the first galaxies by observing their impact on the intergalactic medium (IGM) as they reionized it during the first billion years leads us to ask whether such indirect observations are capable of diagnosing which types of galaxies were most responsible for reionization. We attempt to answer this by extending the galaxy mass range of our first generation of large-scale radiative transfer simulations of reionization in volumes large enough, $\\geq(100/h Mpc)^3$, to make statistically meaningful predictions of observable signatures, downward from those above $10^9\\,M_\\odot$ (high-mass, atomic-cooling halos, or \"HMACHs\") to include those between $10^8$ and $10^9\\,M_\\odot$ (low-mass, atomic-cooling halos, or \"LMACHs\"), as well. Previously, we simulated the effects of both HMACHs and LMACHs but only by reducing the box size to $35/h$~Mpc, too small to apply those simulations to predict observables like the 21-cm background fluctuations.  Those simulations showed that LMACHS can make a difference, however. While LMACHs are even more abundant, photoheating suppresses their ability to form stars if they are located inside ionized regions of the IGM, so their contribution starts reionization earlier but tends to saturate over time before it ends. To predict the observables in that case while explicitly tracking the radiation from this full mass range of galaxies, we have had to advance both our N-body and radiative transfer methods substantially, as described here. With this, we perform new simulations in a box 163~Mpc on a side, and focus here on predictions of the 21cm background, to see if upcoming observations are capable of distinguishing a universe ionized primarily by HMACHS from one in which both HMACHs and LMACHs are responsible, and to see how these results depend upon the uncertain source efficiencies.  We find that 21-cm fluctuation power spectra observed by the first generation EoR/21cm radio interferometer arrays should be able to distinguish the case of reionization by HMACHs alone from that by both HMACHs and LMACHs, together. Some reionization scenarios, e.g. one with abundant low-efficiency sources vs. one with self-regulation, yield very similar power spectra and rms evolution and thus can only be discriminated by their different mean reionization history and 21-cm PDF distributions. We find that the skewness of the 21-cm PDF distribution smoothed over LOFAR-like window shows a clear feature correlated with the rise of the rms due to patchiness. This is independent on the reionization scenario and thus provides a new approach for detecting the rise of large-scale patchiness and an independent check on other measurements, regardless of the detailed properties of the sources. The peak epoch of the 21-cm rms fluctuations depends significantly on the beam and bandwidth smoothing size as well as on the reionization scenario and can occur for ionized fractions as low as 30\\% and as high as 70\\%. Measurements of the mean photoionization rates are sensitive to the average density of the regions being studied and therefore could be strongly skewed in certain cases. Finally, the simulation volume employed has very modest effects on the results during the early and intermediate stages of reionization, but late-time signatures could be significantly affected. ",
        "introduction": "Study of the Epoch of Reionization (EoR) has progressed in recent years in response to a number of new observational developments. The combination of the CMBR data from WMAP \\citep{2011ApJS..192...18K,2011ApJS..192...16L} and ever deeper ground-based observations of high-redshift QSOs, galaxies and GRBs \\citep{2010ApJ...723..869O, 2011ApJ...734..119K,2011Natur.474..616M,2011ApJ...736....7C, 2011arXiv1106.6055K} clearly suggest that the reionization process started early and was quite extended in time. However, observations of the effects of the EoR are just beginning to assemble constraints sufficient to diagnose the conditions which brought it about. Ongoing and upcoming observations are expected to put further, much more stringent constraints on the reionization history. The best constraints are likely to result from redshifted 21-cm experiments with the low-frequency radio interferometers GMRT\\footnote{\\url{http://gmrt.ncra.tifr.res.in/}} \\citep{2011MNRAS.413.1174P}, LOFAR\\footnote{\\url{http://www.lofar.org/}} \\citep[e.g.][]{2010MNRAS.405.2492H}, MWA \\citep{2009IEEEP..97.1497L}\\footnote{\\url{http://www.mwatelescope.org/}} and PAPER \\citep{2010AJ....139.1468P}. Additional information will come from Planck satellite \\citep{2011arXiv1101.2022P} and measurements of the near infrared background with CIBER\\footnote{\\url{http://physics.ucsd.edu/~bkeating/CIBER.html}} and AKARI\\footnote{\\url{http://irsa.ipac.caltech.edu/Missions/akari.html}}, among others.  The understanding and correct interpretation of these observational results requires detailed modelling. Specific characteristics and features of the observable signatures provide information about different aspects of EoR. One of the central questions is what are the nature, abundances and physical properties of the ionizing sources. Our purpose in this work is to explore how observations of the EoR might diagnose the nature of the reionization sources.  Our first generation of simulations of inhomogeneous reionization combined cosmological N-body simulations of galaxy and large-scale structure formation and the intergalactic density and velocity fields with detailed radiative transfer calculations of the ionizing radiation from every galactic halo whose formation we resolved in a volume large enough to make statistically meaningful predictions of observable consequences of reionization \\citep{2006MNRAS.369.1625I, 2006MNRAS.372..679M,2008MNRAS.384..863I,kSZ,cmbpol,2008MNRAS.391...63I, 2009MNRAS.393.1449H,2010MNRAS.406.2521I,2010ApJ...710.1089F}.  These simulations were in comoving boxes of size 100$/h$ ($=143$ for $h=0.7$ henceforth) Mpc on a side. Previous simulations had been limited to much smaller volumes, too small to serve this purpose. Earlier simulations of smaller volumes, for example, underestimated the width of the time interval for the global transition of the intergalactic medium (IGM) from neutral to ionized \\citep{2000ApJ...535..530G,2002ApJ...575...33R,2003MNRAS.344..607S, 2003MNRAS.344L...7C}, as well as the amplitude of the kSZ fluctuations in the temperature of the CMB from the EOR \\citep{2001ApJ...551....3G, 2005MNRAS.360.1063S}. The characteristic size of the intergalactic H~II regions during the EOR is expected to reach 10's of Mpc \\citep{2004ApJ...613....1F,2006MNRAS.365..115F,2011MNRAS.413.1353F} before they grow large enough to overlap. Any fluctuations introduced by this ``patchiness'' scale, therefore, require simulation volumes at least this large to model reionization. Moreover, since the first generation of radio observations seeking to detect fluctuations in the brightness temperature of the 21cm background from the EoR have angular resolution of a few arcminutes, a simulation box size in excess of $\\sim100$~Mpc is necessary to characterize the power spectrum, at wavenumbers as small as $\\sim 0.1\\,{\\rm Mpc}^{-1}$ where the initial sensitivity will peak.  Our simulations of a comoving volume 100$/h$ (= 143) Mpc on a side were limited, however, by the mass resolution of the N-body simulations, to the direct simulation of galactic halo sources more massive than $\\sim 2\\times10^9\\,\\msun$.  Halo sources are also possible at lower mass if halo gas can radiatively cool below the halo virial temperature to make star formation possible.  This includes halos above about $10^8\\,\\msun$, for which collisional excitation of H atoms can radiatively cool the primordial-composition halo gas since the virial temperature is above $10^4$ K.  We shall refer to these halos between about $10^8$ and $10^9\\,\\msun$ as low-mass atomic-cooling halos (``LMACHs''), to distinguish them from the halos above $10^9\\,\\msun$, which we shall call high-mass, atomic-cooling halos (``HMACHs''). LMACHs are more numerous than HMACHs at these epochs, so it might be thought that they would dominate reionization. However, unlike the HMACHs, the LMACHs are vulnerable to the negative feedback effects of photo-heating if they form inside a pre-exisiting H~II region of the IGM, since the pressure of the IGM would then prevent the intergalactic gas from collapsing gravitationally into their dark matter host halos \\citep{1992MNRAS.256P..43E,1994ApJ...427...25S,1996MNRAS.278L..49Q, 1997ApJ...478...13N,2004ApJ...610L...5S,2008MNRAS.390..920O, 2008MNRAS.390.1071M}.  As first emphasized by \\citet{1994ApJ...427...25S}, this limits their contribution to reionization. Halos of even smaller mass than LMACHs would be even more vulnerable to the negative feedback effects of photoionization heating, since they would, in addition to being prevented from capturing intergalactic gas, photoevaporate whatever interstellar gas they had already accumulated before they were engulfed by reionization \\citep{2004MNRAS.348..753S,2005MNRAS...361..405I}.  Before that reionization, however, these minihalos with masses below about $10^8\\,\\msun$, nevertheless could have formed stars if enough H$_2$ molecules were present in them.  Since their virial temperatures are below $10^4$ K, their gas is too cold for collisional excitation and radiative cooling by H atoms, but cooling through the collisionally exciting rotational-vibrational levels of H$_2$ is possible. However, these molecules were easily dissociated by the rising UV background of starlight at energies below the ionization threshold of H atoms, an inevitable by-product of the same stars that contributed to reionization. This tends to limit the contribution of minihalos to reionization early in the EoR \\citep{2000ApJ...534...11H, 2009ApJ...695.1430A}. Previous estimates of the minihalo contribution, as a result, assume that this contribution is smaller than that of the LMACHs and HMACHS, so for now, we shall neglect it, although patchiness in the UV background may make them important than one might naively think \\citep{2009ApJ...695.1430A}.  To investigate the impact of the LMACHs on reionization and its observable properties by direct radiative transfer simulation of reionization, our first generation of simulations boosted the halo mass resolution in order to resolve all the halos of mass $10^8\\,\\msun$ and above, but sacrificed volume by simulating in a box of size 37$/h$ (= 53 Mpc) on a side \\citep{2007MNRAS.376..534I}. These simulations demonstrated explicitly that reionization in the presence of the LMACHs and their suppression if they formed inside pre-existing H~II regions during the EoR was ``self-regulated''. The more LMACHs that formed, the more volume and mass of the IGM was ionized, but as this ionized fraction grew, so did the fraction of the total LMACH halo population that formed inside the ionized regions and was suppressed as sources of reionization.   This meant that, although the LMACHs dominated the early phase of reionization, their contribution to reionization eventually saturated, and reionization was finished by the HMACHs, whose abundance rose exponentially over time.  An observational consequence of importance that is a possible signature of this process was an early onset but late finish for the EoR, as required to explain the high values of electron scattering optical depth reported from WMAP observations of the large-angle fluctuations in the polarization of the CMB. To make further predictions of observable consequences, however, we need to enlarge the volume of these simulations to greater than 100/h Mpc on a side, without losing this enhanced mass resolution necessary to resolve the LMACHs. That is the purpose of the new developments we report in this paper.  To accomplish this goal, it was first necessary to improve and advance both our N-body and radiative transfer methods in order to be able to simulate halo formation with much higher mass resolution and to transfer the ionizing radiation from a much larger number of sources. Toward this end, both codes had to become massively parallel, running on thousands of computing cores, as well as more efficient. These numerical developments are discussed in more detail in \\S~\\ref{sim:sect} below and in \\citet{2008arXiv0806.2887I}.  While the work described here follows naturally from our own previous work as described above, it also differs substantially from other work in the literature to-date involving large-scale radiative transfer simulations of reionization. A full account of that literature is well-beyond the scope of this paper, but we will mention a few points to distinguish the current work. Our N-body simulations resolve all galactic halo sources of $10^8$ $\\msun$ and above, in a comoving box as large as $114/h=163$~Mpc on a side. We post-process the density field of the IGM and the galactic halo source populations derived from these N-body simulations by performing a detailed, ray-tracing calculation on a grid of $256^3$ cells.  The simulations described in \\citet{2007MNRAS.377.1043M} were based upon post-processing N-body simulations with halo mass resolution above $10^9$ $\\msun$, in a box $65.6/h$ Mpc on a side, a volume which is five times smaller than ours, also on a grid of $256^3$ cells. \\citet{2007MNRAS.377.1043M}  were, thus, unable to treat explicitly the halo mass range below $10^9$ $\\msun$, which is subject to suppression by the negative feedback effects of reionization, but they did include a semi-analytical, ``subgrid'' approximation for the contribution from the unresolved, smaller-mass halos, including some feedback effects. \\citet{2007ApJ...671....1T} considered a $50/h$ Mpc box, hence, an order of magnitude smaller volume than ours, and $180^3$ cells for radiative transfer, but their halo mass resolution was similar to ours.  They did not consider the feedback effects of reionization on the small-mass galactic halo sources as we do here. \\citet{2008ApJ...681..756S} subsequently applied this method to simulate a volume $100/h$ Mpc on a side ($2/3$ of our volume), with similar halo mass resolution, on a radiative transfer grid of $360^3$ cells, to study the structure of the patchy ionization field during the EOR, again without considering feedback and doing only a single simulation, without considering any variation of the (highly uncertain) reionization parameters. Recently, the codes of \\citet{2007MNRAS.377.1043M} and \\citet{2007ApJ...671....1T} were compared with each other in \\citet{2011MNRAS.414..727Z}, by comparing results for the ionization fields and 21-cm brightness temperature fluctuation statistics, for reionizaton simulations advanced part-way through the epoch of reionization, to the 72\\% ionized point, based upon post-processing a previously-simulated input density field in a box $100/h$ Mpc on a side, smoothed to a radiative transfer grid with $256^3$ cells, with halo mass resolution of $10^8$ solar masses, without any feedback or back-reaction on the halo sources. They also considered a single reionization scenario, with no parameter variation. Finally, \\citet{2010ApJ...724..244A} simulated radiative transfer in several different boxes, as large as $100/h$ Mpc, by post-processing a density field and galaxy population derived from separate simulations that combined N-body dynamics and hydrodynamics on a $1024^3$ grid, but with minimum resolved galaxy masses in that case as large as $8\\times 10^9$ $\\msun$. No feedback from reionization on galactic sources was considered. While there are other distinctions of detail both between these other simulations and ours and of one from another, we have listed these above to make it clear that our current paper will describe simulations and results which are new, {\\it both} because they are based upon different methodology, applied in greater depth than previously to predict observables like the 21cm background from the EOR, {\\it and} because they are on an unprecedented scale.  The rest of paper is organized as follows. In \\S~\\ref{sim:sect} we present our codes, numerical methods and simulations. In \\S~\\ref{results:sect} we discuss our results on the formation of early cosmic structures. In \\S~\\ref{basic_results:sect} we present our results on the basic reionization features, reionization history, integrated electron-scattering optical depth and geometry of ionized patches. The observational signatures derived from our simulations are presented and discussed in \\S~\\ref{observ:sect}. Our conclusions are summarized in \\S~\\ref{summary:sect}. Finally, in Appendix~\\ref{appendixA} we present the (physically less realistic) cases of reionization by rare, massive sources, while in Appendix~\\ref{appendixB} we discuss the more technical point of the dependence of our results on the Jeans suppression threshold for low-mass sources. ",
        "conclusions": "\\label{summary:sect}  In this work we have used a large set of cosmological structure formation and reionization simulations in an attempt to gain insight into what can be learned about the properties of the reionization sources based on their observational signatures. In particular, we are interested in determining what observations could be used to discriminate between certain source models, thereby restricting the available parameter space. Here we primarily focused on the redshifted 21-cm signatures, as these can in principle probe the full reionization history and offer a wide range of different probes, from the mean history, through fluctuations measures like rms evolution and power spectra, to PDFs and higher-order statistics which can detect non-Gaussian features.  The observable features of the epoch of reionization derive from the gradual mean transition of the IGM from neutral to highly ionized state, as well as from the patchiness of that transition. The mean transition depends largely on the overall number of ionizing photons emitted by the source per unit time, with some correction due to recombinations which is position-dependent due to density spatial variations. The patchiness, on the other hand depends, in a complicated way on the abundances, clustering and efficiencies of the ionizing sources.  Our structure formation simulations confirm previous results by us and other groups that the high-redshift halo mass functions are inconsistent with either of the widely-used Press-Schechter and Sheth-Tormen analytical fits. In particular, the abundance of rare halos is strongly underestimated by PS, but over-estimated by ST. We find that the nonlinear halo bias is extremely high and very scale-dependent. Linear bias regime is only reached at very large scales, $k\\gtrsim0.1$. For the rarest halos (3-$\\sigma$ and above) linear bias regime is never reached even within our largest, 163~Mpc volume. Therefore, proper account for the nonlinear bias of halos is important and any calculations assuming linear bias are underestimating the halo clustering significantly.  The Jeans mass filtering of low-mass halos in ionized regions and the related self-regulation of reionization results in significantly more extended reionization history and higher integrated electron scattering optical depth (by $\\Delta\\tau_{\\rm es}\\sim0.01$) compared to the high-mass source-only scenario with the same overlap redshift, albeit both optical depths are still within the current constraints from WMAP. Even more significant are the changes in the reionization geometry, resulting in corresponding differences in the reionization observables. The 21-cm fluctuations are lower at all scales and their PDF distributions are somewhat more Gaussian, although significant non-Gaussianity remains.  In all our simulations reionization occurs inside-out, with the high-density regions being reionized on average earlier than the mean and low-density ones. This inside-out nature of reionization results in the mass-weighted IGM photoionization rates being considerably larger, by factor of a few, than volume-weighted ones. This should always be taken into account as it can easily skew observational measurements depending on the mean density of the regions being probed.  The skewness of the 21-cm PDF distribution smoothed over LOFAR-like window shows a clear feature correlated with the rise of the rms due to patchiness. This feature does not exist in the unsmoothed data, indicating that it is related to the non-Gaussianity of the large-scale patchy distribution of 21cm emission. The feature exists for any reionization scenario and ionizing source properties and thereby provides a different approach for detecting the rise of large-scale patchiness and an independent check on other detections.  The peak position of the 21-cm rms fluctuations depends significantly on the beam and bandwidth smoothing size as well as on the reionization scenario. As a consequence, it does not always occur at 50\\% ionization fraction as sometimes is claimed, but instead can happen for ionized fractions as low as 30\\% and as high as 70\\%.  The simulation volume has only a modest effect on the results as long as the typical size of the ionized patches is smaller than the volume. However, the fluctuations at large scales (above approximately a fifth of the boxsize) are severely affected. This is especially important at late times, when the ionized patches grow very large. Therefore, at least $\\sim100/h$~Mpc boxes are required to model the fluctuations during the late stages of reionization.  The ionizing source efficiencies and their correlation properties introduce clear signatures in the reionization observables. As a direct consequence of that, one cannot model low-mass, unresolved sources by simply assigning their emissivity to the resolved higher-mass sources as the latter have abundance which is highly variable over time and different clustering properties from lower-mass sources which provide the bulk of the ionizing photons (see Appendix~\\ref{appendixA}).  When self-regulation is present there are only minor differences between the 21-cm observational signatures resulting from high- and low-efficiency ionizing sources, apart from an overall shift of the reionization history. The corresponding PDF distributions are also very similar, which suggests that the source efficiencies in such models can only be constrained by the overall timing of the mean reionization history.  Scenarios where low-mass sources are completely absent, e.g. somehow rendered sterile, are relatively easily distinguishable from the ones where they are present (even if strongly suppressed). On the other hand, our results suggest that numerous low-efficiency sources (case S4) can mimic the effects of suppression (S1). Such scenarios therefore might be difficult to distinguish solely based on power spectra and similar measurements. However, they might still be discriminated through 21-cm PDFs as the no suppression case creates many more high-brightness peaks. Similarly, a high-mass source only scenario (S5) gives quite similar results to the high-efficiency, no suppression case (S3) at the same stages of reionization (albeit these cases overlap at different times for the parameters we have chosen), although they do differ in terms of power at large scales and in number of bright peaks.  The results presented in this work should not be considered to be an ultimate, realistic prediction of the reionization signals. While the assumptions about source efficiencies and suppression we made for our fiducial cases are reasonable based on our current knowledge and likely bracket the realistic range, the uncertainties are still substantial. As more observational data becomes available over time it can be used to restrict the parameter space further and help us refine our theoretical models, which in turn will provide a valuable tool for interpreting the meaning of the observational results in terms of early structure formation, source efficiencies, suppression mechanisms, etc. In this framework our current study, which evaluates in a controlled way the effects of a set of widely different assumptions about the sources of ionizing radiation is a very useful step towards a more complete understanding of early galaxy formation and feedback."
    },
    "1107/1107.1682_arXiv.txt": {
        "abstract": "We present here the All Sky Transmission MONitor (ASTMON), designed to perform a continuous monitoring of the surface brightness of the complete night-sky in several bands. The data acquired are used to derive, in addition, a subsequent map of the multiband atmospheric extinction at any location in the sky, and a map of the cloud coverage. The instrument has been manufactured to afford extreme weather conditions, and remain operative. Designed to be fully robotic, it is ideal to be installed outdoors, as a permanent monitoring station. The preliminary results based on two of the currently operative units (at Do\\~nana National Park - Huelva- and at the Calar Alto Observatory - Almer\\'\\i a -, in Spain), are presented here. The parameters derived using ASTMON are in good agreement with previously reported ones, what illustrates the validity of the design and the accuracy of the manufacturing.  The information provided by this instrument will be presented in forthcoming articles, once we have accumulated a statistically amount of data. ",
        "introduction": "The night sky brightness, the number of clear nights, the seeing, transparency and photometric stability are some of the most important parameters that qualify a site for front-line ground-based astronomy (Taylor et al. 2004). There is limited control over all these parameters, and only in the case of the sky brightness is it possible to keep it at its natural level by preventing light pollution from the immediate vicinity of the observatory (Garstang 1989). Previous to the installation of any observatory, extensive tests of these parameters are carried out in order to find the best locations, maximizing then the efficiency of these expensive infrastructures (Thomas-Osip et al. 2010, Schock et al. 2009) . However, most of these parameters are not constant along the time, neither in short term nor in long term time scales.  An on-line monitoring of all these conditions has been proved to be a fundamental tool to take decisions on the observational strategies, in order to increase the efficiency of any professional observatory (Travouillon et al. 2011). In particular, a very interesting input to decide about the optimal strategy would be to estimate in real-time the stability of the atmosphere and the location of the darkest and the most transparent areas in the sky (eg, Benn \\& Ellison et al. 1998a,b; S\\'anchez et al. 2007; Moles et al. 2010; High et al. 2010; Zou et al. 2010).   Along the 26 years of operations of the Calar Alto observatory, there has been different attempts to characterize some of the main properties described before: (i) Leinert et al. (1995) determined the sky brightness corresponding to the year 1990; (ii) Hopp \\& Fernandez (2002) studied the extinction curve corresponding to the years 1986-2000; However, we lack of a consistent study of all of them, spanning over a similar time period.This information becomes even more important for long-term surveys, where automatic or semi-automatic schedulers are programmed in order to optimize the observation over large areas in the sky (eg., Moles et al. 2008).  In order to characterize these parameters most of the professional observatories have installed permanent monitors which provide an almost real-time estimation of them. The values produced by these monitors are frequently accessible on-line via web services, stored in databases, and even included in the headers of the images obtained at the observatory. They compose an extremely useful database on the long term stability of the night-sky conditions at a certain site. In general, most of these monitors produce global parameters, based on large aperture measurements, like the cloud monitors based on the temperature difference between the ground layer and the sky, or local parameters, like the information provided by most of the seeing and transparency monitors: eg., RoboDIMM (Aceituno 2004) or EXCALIBUR (S\\'anchez et al. 2007, Moles et al.  2010). On other occasions, all sky monitors provide qualitative estimations on the sky transparency, based on direct images to estimate the cloud coverage, but no quantitative ones.  As part of a large program to determine the observing conditions at different locations in Spain (S\\'anchez et al. 2007, 2008; Moles et al. 2010), we developed a new set of instruments, able to estimate the most important parameters that determine the quality of an astronomical site, working in fully automatic mode. The first of these instruments was named EXCALIBUR (P\\'erez-R\\'amirez et  al. 2008), a multi-wavelength extinction monitor, whose main purpose was to determine the atmospheric attenuation curve, and derive the relative contribution of each its components, at a particular location in the sky. The main idea behind this instrument was to determine these parameters at the sample position where the astronomical observation is taking place, monitoring any possible change in the night-sky stability.  This instrument provides accurate local information, but it cannot provide simultaneous information on the considered parameters at different locations in the sky. For doing so, a different concept of instrument is required. In this article we present the All Sky Transmission MONitor (ASTMON), an instrument that, working on the same principals of EXCALIBUR, derives the surface sky brightness, atmospheric extinction and extinction curve at several locations in the sky simultaneously. It also provides additional information about the night-sky quality, based on the measurements acquired, like the photometric conditions, cloud-coverage and amount of light-pollution.  Although the main goal of these instruments is to determine the night-sky quality for the astronomical observation, this information has been found useful in another science fields. Due to that, a prototype of a EXCALIBUR unit was installed in the Andalusian Center of Environment Studies (Granada), to estimate the light pollution and aerosol content in populated areas (P\\'erez-Ram\\'\\i rez et al. 2008). In addition, the first prototype of ASTMON was installed permanently in the Do\\~nana Biological Reserve, CSIC (Spain), to determine the effects of the light pollution in the biological environment of this protected area. Both instruments were calibrated at the Calar Alto observatory, which has acquired and installed units of both of them, integrating them in their monitor systems. The calibration runs demonstrated that these instruments are able to provide quantitative valuable measurements in a completely automatic way, comparable to the values provided by other monitors and/or guided astronomical observations.  The structure of this article is as follows: in Section \\ref{con} we present the basic concepts in which the instrument is based; in Section \\ref{inst} the instrument itself is described, illustrating its main components; in Section \\ref{data} we present the results from the calibration run, comparing them with those derived by other well tested monitors and/or published data. In Section \\ref{sum}, we present a summary of the results from the experiment and the future perspective of this kind of instruments. ",
        "conclusions": "\\label{sum}  We have developed and manufactured the first permanent station to derive an exhaustive study of the light pollution and the atmospheric extinction, by applying the standard techniques of astronomical photometry. The system is based on a all-sky fisheye lens, with a filter wheel, and a CCD detector that allows to obtain automatic measurements of the night-sky surface brightness in 5 different bands.  The system has been designed to afford extreme weather conditions and it can be installed in a remote place, being fully robotic. The instrument comprises a software that it is fully automatic. It performs all the required tools to operate the instrument, switching on and off automatically, acquires the science frames in the corresponding filters, and finally, reduces, calibrates and analyzes them to derive time dependent maps of the night-sky surface brightness, and individual estimations of atmospheric extinction for every band, together with a rough estimation of the cloud coverage.  The early experiments performed during the calibration runs of the instrument have demonstrated the validity of the conceptual design and its implementation. The values provided by the instrument, for both the night-sky surface brightness at the zenith, and the dust extinction, are consistent with ones reported in previous published results.  Currently, three ASTMON units have been installed at different locations in Spain: (i) the Calar Alto Observatory (Almeria); (ii) the Do\\~nana National Park (Huelva) and (iii) the building of the Physics Department of the Universidad Complutense de Madrid (Madrid). All of them are fully operational and they are providing data for every clear night. A deep analysis of these data will be presented in forthcoming articles, once we have acquired a statistically significant number of them."
    },
    "1107/1107.5297_arXiv.txt": {
        "abstract": "High spectral-resolution Br\\,$\\gamma$ VLTI/{\\sc AMBER} observations of classical Be stars reveal complex visibility, phase, and closure phase profiles, which were not seen in previous lower-resolution data.  They present new challenges for the modeling of Be stars and may require additional ingredients over and above the conventional central star + circumstellar disk paradigm. ",
        "introduction": "{\\sc AMBER} \\citep{2007A&A...464....1P} was the first high-spectral resolution (HR, R~$=$~12\\,000) interferometric instrument and still is the only one capable of taking closure phase data in this mode.  After the commissioning in the end of 2008, the spectral resolution was increased eightfold wrt. its medium-resolution mode, and it is about 30 times as high as  that of CHARA/{\\sc MIRC} \\citep{2004SPIE.5491.1370M}. The present (2011) list of standard HR settings comprises 14 central wavelengths but most of the actual HR observations cover the Br~$\\gamma$ spectral region.  The order-of-magnitude increase in spectral resolution entails the potential for a qualitative leap forward in studies of objects with a considerable range in dynamics in general and rotationally supported disks of young or evolved stars in particular, which account for a large fraction of all available VLTI observations.  In order to explore this potential for classical Be stars, we have combined data from three {\\sc AMBER} projects of our own with selected archival observations.  The objects in Tab.~\\ref{targets} include almost all Be stars observed to date with {\\sc AMBER} in its HR mode. Only Achernar ($\\alpha$\\,Eri) was not included because it had but a tiny disk at the time of the {\\sc AMBER} observations. Observations were done both with the 8.2m diameter Unit Telescopes (UTs) and the 1.8m Auxiliary Telescopes (ATs). Because even the largest circumstellar Be disk diameters in Br\\,$\\gamma$ are of a few miliarcseconds, the longest baselines -- UT1-UT3-UT4 or UT1-UT2-UT4 triplets, D0-H0-G1-I1 or A0-K0-G1-I1 ATs quadruplets -- were used for observations. All data were reduced with the help of the yorick amdlib3 reduction package \\citep{2007A&A...464...29T,2009A&A...502..705C}.   \\begin{table}[!ht] \\caption{Classical Be stars observed in {\\sc AMBER} HR mode and Br\\,$\\gamma$ region in 2008-2010. Parameters of the targets are taken from the Simbad database.} \\smallskip \\begin{center} {\\small \\begin{tabular}{rrlrrrl} \\tableline \\noalign{\\smallskip} Star             &  HD    & Spectral & v $\\sin i$ & $K_{\\mathrm{mag}}$   & parallax  & Epoch          \\\\ &        &  type    & [km/s]     & [mag]             &  [mas]    &                \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} $\\zeta$\\,Tau     &  37202 & B2IV     &  310       &         2.81      &  7.82     &  Oct 15, 2008  \\\\ $\\alpha$\\,Col    &  37795 & B7IVe    &  195       &         2.83      & 12.16     &  Jan 9,10, 20, 2010  \\\\ $\\kappa$\\,CMa    &  50013 & B1.5IVe  &  210       &         3.55      &  4.13     &  Jan 9, 10, 17, 2010   \\\\ $\\omega$\\,CMa    &  56139 & B2IV-Ve  &  105       &         4.37      &  3.53     &  2008-2010, large dataset  \\\\ $\\beta$\\,CMi     &  58715 & B8Ve     &  210       &         3.10      & 19.16     &  Dec 31, 2009  \\\\ $\\omega$\\,Car    &  89080 & B8IIIe   &  220       &         3.45      &  8.81     &  Dec 21, 2008 \\\\ 48\\,Lib          & 142983 & B5IIIp(?)&  395       &         4.59      &  6.36     &  May 12, 2009 \\\\ $\\delta$\\,Sco    & 143275 & B0.2IVe  &  165       &         2.43      &  8.12     &  2009-2011; continuing obs. \\\\ \\noalign{\\smallskip} \\tableline \\label{targets} \\end{tabular} } \\end{center} \\end{table}  \\begin{figure}[!ht] \\includegraphics[width=\\textwidth,clip=true]{NMpkCMa.ps} \\caption{$\\kappa$\\,CMa observed on January 9, 2010, AT baseline configuration D0-H0-K0. The baseline lengths and position angles are: D0-H0 64.0m@71.0$^{\\circ}$, H0-K0 32.0m@71.0$^{\\circ}$ and D0-K0 96.0m@71.0$^{\\circ}$. Although canonical S-shaped profiles appear in phase diagrams,  a double-peak reversal can be seen in all visibilities.} \\label{kCMaint} \\end{figure} ",
        "conclusions": "The {\\sc AMBER} HR database of classical Be stars documents that both double/multiple visibility profiles and phase reversals are common features.  Double/multiple-peak visibility profiles appeared in all observed Be stars at some baseline(s) and time(s).  Their occurrence and peak separation do not seem to depend on the length of the baseline.  Asymmetric visibility profiles may correspond to a transition between single and double peak profiles.  Because not even in the most nearby objects in our sample the central stars are resolved, the double-peak visibility profiles are not caused by the superposition of stellar and disk visibilities.  Phase reversals occurred in five of the eight observed Be stars: ($\\zeta$\\,Tau, $\\alpha$\\,Col, $\\beta$\\,CMi,  48\\,Lib, $\\delta$\\,Sco, and $\\omega$\\,CMa). The preferential detection of phase reversals at long baselines may locate the responsible physical process close to the star. In $\\beta$\\,CMi, the feature is strong also at the 45m UT1-UT2 baseline; however, this is a very nearby star.  Two of the stars showing a phase reversal are Be shell stars seen equator-on ($\\zeta$\\,Tau, 48\\,Lib). However, the limited database (only one  HR observation for 4 targets, 3 observations of different quality for 2 targets) does not permit phase reversals and visibilities with more than one peak to be linked to disk position angle. For 28\\,CMa and $\\delta$\\,Sco this point will be discussed in dedicated papers. Nevertheless, the small sample indicates that these features occur over a broad range of spectral B sub-types and disk viewing angles, from equator-on to at least medium inclination.  Their presence in steady-state phases of the disks around three stars ($\\alpha$\\,Col, $\\beta$\\,CMi, $\\kappa$\\,CMa) and also shortly after an outburst in $\\omega$\\,CMa suggests that the observed phase reversals are unrelated to a special disk evolution phase.  The implied commonality in Be star disks of multiple-peak visibility profiles and phase reversals represents a major challenge for disk models. They do not match the S-shaped disk phase profiles, which are observed at lower spectral resolution and form the basis of commonly accepted models. Without quantitative modeling, their interpretation is rather speculative.  Our preliminary efforts (Carciofi et al., in preparation) suggest that secondary dynamical disk effects may explain double-peak visibility profiles and phase reversals at baselines perpendicular to the assumed equatorial plane (48\\,Lib, $\\delta$\\,Sco). This is crudely consistent with the observed position angle-dependence in some stars.  However, the model is struggling with phase reversals at baselines parallel to the equatorial plane (as in $\\alpha$\\,Col or $\\beta$\\,CMi).  Therefore, an alternative view is developed in Rivinius et al.\\ (A\\&A, to be submitted).  With polar mass outflows, it goes beyond the canonical central star $+$ circumstellar disk paradigm for classical Be stars but may not be applicable to all classical Be stars.  Obviously, high spectral resolution adds a new and essential dimension to interferometry of objects with significant internal motions.  This is demonstrated by the here presented observations of classical Be stars."
    },
    "1107/1107.2648_arXiv.txt": {
        "abstract": "Rapid mass assembly, likely from mergers or smooth accretion, has been predicted to play a vital role in star-formation in high-redshift \\lya\\ emitters. Here we predict the major merger, minor merger, and smooth accreting \\lya\\ emitter fraction from $z \\approx3$ to $z\\approx7$ using a large dark matter simulation, and a simple physical model that is successful in reproducing many observations over this large  redshift range. The central tenet of this model, different from many of the earlier models, is that the star-formation in \\lya\\ emitters is proportional to the mass accretion rate rather than the total halo mass. We find that at  $z\\approx3$,  nearly $35\\%$ of the \\lya\\  emitters accrete their mass through major (3:1) mergers, and this fraction increases to about $50\\%$ at $z\\approx 7$. This imply that the star-formation in a large fraction of high-redshift \\lya\\ emitters is driven by mergers. While there is discrepancy between the model predictions and observed merger fractions, some of this difference ($\\sim15\\%$)  can be attributed to the mass-ratio used to define a merger in the simulation. We predict that future, deeper observations which use a 3:1 definition of major mergers will find $\\geq 30\\% $ major merger fraction of \\lya\\ emitters at redshifts $\\geq 3$. ",
        "introduction": "While several theoretical models of \\lya\\ emitters have been developed \\citep[e.g.][]{bar04,dav06,tas06,shi07,nag08,del06,kob07,kob09,day08,sam09} we still lack a complete understanding of how mass assembly  and star-formation occurs in these galaxies. It is likely that mergers are the dominant mode of mass accretion resulting in star-formation in at least some fraction ($\\geq20\\%$)  of \\lya\\ emitters  \\citep[e.g.][]{pir07,bon09,tan09}.   Recently,  \\citet{til09} developed a physical model of \\lya\\ emitters which is successful in explaining many observable including number density, stellar mass, star-formation rate, and clustering properties of \\lya\\ emitters from $z\\approx 3$ to $z\\approx 7$. This model differs fundamentally from many of the earlier models \\citep[e.g.][]{hai99,dij07,mao07,sta07,fer08} in that the star-formation of \\lya\\  emitters, and hence their \\lya\\ luminosity is proportional to the mass accretion rate rather than the total halo mass.   In this paper, we combine the above model with a large dark matter cosmological simulation to predict the major merger, minor merger, and smooth accreting \\lya\\ emitter fraction from $z\\approx3$ to $z\\approx 7$. We also carefully asses the uncertainties in our model predictions. We note that currently there are no theoretical prediction of merger fraction of \\lya\\ emitters.   On the observational front,  only recently  it  has been possible  to estimate the merger fraction of \\lya\\ emitters. At $z\\sim 0.3$,  \\citet{cow10} found that $\\geq30\\%$   of \\lya\\ emitters are either irregulars or have disturbed morphologies indicative of mergers. At higher redshifts, $z > 3$, the observed merger fraction varies from $\\approx 20\\%$ to $\\approx 45 \\%$ \\citep[]{pir07,bon09,tan09}.  Recently, \\citet{coo10} found that all Lyman-break galaxies (LBGs)  that have  close companions ($\\leq 15 h^{-1}$ Mpc) exhibit  strong \\lya\\ emission lines. In the spectroscopically confirmed close pairs, the spectra/imaging show double \\lya\\ emission peak/double morphology confirming that these LBGs at $z\\approx3$ are indeed mergers. \\citet{coo10} also studied a sample of spectroscopically confirmed LBGs at $z\\approx3$ from \\citet{sha06}, and found that about 33\\% of LBGs with strong \\lya\\ emission have double \\lya\\ emission peaks, strengthening the above conclusion.    The outline of this paper is as follows. In \\S2 we briefly describe the physical model of \\lya\\ emitters \\citep{til09}, our  simulations,  and how we construct the merger trees. In  \\S3,  we estimate the merger fraction of \\lya\\ emitters, discuss the effect of  mass resolution in the simulation on our results, and compare our model predictions with the observations. In \\S4 we quantify the  uncertainties associated with the predictions, and we summarize our results  in \\S5. ",
        "conclusions": "Motivated by our earlier work \\citep{til09} based on a simple physical model of \\lya\\ emitters, combined  with a large dark matter cosmological simulation, we have predicted the major merger, minor merger, and smooth accreting \\lya\\ emitter fraction from $z\\approx 3$ to $z\\approx7$. The model presented in \\citet{til09}, different from many of the earlier models in that the star-formation rate is proportional to the mass accretion rate rather than the total halo mass, has been  successful in reproducing many  observed physical properties including the luminosity functions, stellar ages, star formation rates, stellar masses, and clustering properties of \\lya\\ emitters from $z\\approx3$ to $z\\approx7$.    We carefully constructed the merger tree from $z\\approx7$ to $z\\approx3$, accounting for the uncertainties in the halo finder. In this merger tree, each descendent halo was associated with its progenitor(s), and based on the progenitor mass ratio, each merger was classified as either major, minor merger or smooth accretion. We also carefully assessed  how our predicted results might be affected by several parameters including   halo  mass resolution,  defining merger mass ratio, stellar ages of \\lya\\ emitters, and the merger observability timescale. We summarize our key results below:  \\begin{itemize} \\item At $z\\approx3$, about $35\\%$ of \\lya\\ emitters accrete their mass through major mergers, and this  fraction increases to about 50\\% at redshift $z\\approx7$. On the other hand, at this redshift, the minor mergers and smooth accretion contribute equally in the remaining 50\\% of the \\lya\\ emitters. Thus, at higher redshifts, mergers play an important role in mass assembly and hence the star-formation in majority of the \\lya\\ emitters. \\item We also compared our model predicted major merger fraction with the observations and found that our model  over-predicts the major merger fraction. This discrepancy can be resolved if we take into account the uncertainties in defining the merger classification mass ratio. By changing the major merger mass ratio from 3:1 to 2:1, we found that the predicted major merger fraction decreases by about $15\\%$. We predict that future, deeper imaging surveys and using spectroscopic methods such as the one demonstrated in \\citet{coo10}, should find  $\\geq 30\\%$ merger fraction of \\lya\\ emitters at high-redshifts. \\end{itemize}   This work was supported in part by the  grant HST-0808165. All simulations were performed on the $saguaro$ cluster operated by the Fulton School of Engineering at Arizona State University."
    },
    "1107/1107.2362_arXiv.txt": {
        "abstract": "The velocity pattern of a fan loop structure within a solar  active region over the temperature range 0.15--1.5~MK is derived using data from the EUV Imaging Spectrometer (EIS) on board the \\hinode\\ satellite. The loop is aligned towards the observer's line-of-sight and shows downflows (redshifts) of around 15~\\kms\\ up to a temperature of 0.8~MK, but for temperatures of 1.0~MK and above the measured velocity shifts are consistent with no net flow. This velocity result applies over a projected spatial distance of 9~Mm and demonstrates that the cooler, redshifted plasma is physically disconnected from the hotter, stationary plasma. A scenario in which the fan loops consist of at least two groups of ``strands'' -- one cooler and downflowing, the other hotter and stationary -- is suggested. The cooler strands may represent a later evolutionary stage of the hotter strands. A density diagnostic of \\ion{Mg}{vii} was used to show that the electron density at around 0.8~MK falls from $3.2\\times 10^9$~cm$^{-3}$ at the loop base, to $5.0\\times 10^8$~cm$^{-3}$ at a projected height of 15~Mm. A filling factor of 0.2 is found at temperatures close to the formation temperature of \\ion{Mg}{vii} (0.8~MK), confirming that the cooler, downflowing plasma occupies only a fraction of the apparent loop volume. The fan loop is rooted within a so-called ``outflow region'' that displays low intensity and blueshifts of up to 25~\\kms\\ in \\ion{Fe}{xii} \\lam195.12 (formed at 1.5~MK), in contrast to the loop's redshifts of 15~\\kms\\ at 0.8~MK. A new technique for obtaining an absolute wavelength calibration for the EIS instrument is presented and an instrumental effect, possibly related to a distorted point spread function, that affects velocity measurements is identified. ",
        "introduction": "Amongst the complex and varied range of plasma structures in the solar corona, active region fan loops stand out as large, long-lived structures.  The loops are most distinctive when observed in emission lines formed in the narrow temperature range 0.6--1.0~MK, and they came to prominence when the Transition Region and Coronal Explorer (TRACE) satellite began regular observations in an extreme ultraviolet filter centered at 174~\\AA\\ \\citep{handy99}, picking up strong emission lines of \\ion{Fe}{ix} and \\ion{Fe}{x} that are formed at 0.8 and 1.0~MK, respectively. This filter is referred to as ``TRACE 171'' to be consistent with the terminology used for the earlier EUV Imaging Telescope \\citep[EIT;][]{delab95} instrument on board the Solar and Heliospheric Observatory \\citep[SOHO;][]{domingo95} although  in fact the filter is more sensitive to \\ion{Fe}{x} \\lam174.53 \\citep{handy99}. \\citet{schrijver99} presented properties of the active region fan loops seen in the TRACE 171 filter which we summarize here. They are seen to terminate in the penumbra of sunspots and also in strong fields at the edges of active regions; they show small temperature variation with height; they can survive for hours to days; and they show propagating intensity perturbations (PIPs) that give the impression of upflows from the surface. The latter property has been extensively studied, and the most common interpretation is that PIPs are due to upward-propagating, slow magnetoacoustic waves \\citep[e.g.,][]{demoortel02,wang09}. Note that the PIPs are not a permanent feature of fan loops but are intermittent, lasting for typically 10's of minutes before switching off \\citep{mcewan06}. This work also demonstrated that individual `strands' within the loop can display PIPs independently.  Prior to TRACE, the presence of spiky, plumelike structures at the edges of active regions had been remarked on in analyzes of \\emph{Skylab} data. For example, \\citet{foukal76} identified \\ion{Ne}{vii} ($T=0.6$~MK) plumes close to a sunspot, while \\citet{cheng80} found emission ``spikes'' in \\ion{Ne}{vii} at the periphery of an active region, with \\ion{Mg}{ix} ($T=1.0$~MK) loops corresponding with \\ion{Ne}{vii} although the emission was more extended and less sharply defined. These structures are almost certainly the same as the TRACE fan loops.  The first line-of-sight (LOS) velocity measurements for a fan loop were made by \\citet{winebarger02} using the SUMER \\citep[Solar Ultraviolet Measurements of Emitted Radiation;][]{wilhelm95} instrument on board SOHO. A TRACE 171 fan loop was clearly identified in the \\ion{Ne}{viii} \\lam770 emission line (0.7~MK) observed by SUMER and redshifts of 15--40~\\kms\\ were measured, i.e., plasma is downflowing in the fan loop legs. This is in striking contrast to the apparent upflows in the TRACE image sequences first described by \\citet{schrijver99}, and appears to confirm the standard interpretation of the PIPs as due to magnetoacoustic waves rather than an upward flow of plasma.  This view has been challenged more recently by \\citet{depontieu10} who interpret the PIPs as quasi-periodic bursts of upflowing plasma with velocities of 50--150~\\kms. The emission from this plasma is weak compared to that from bulk loop plasma, but is sufficient to give an intensity oscillation signal (which can be detected with an imaging instrument), coupled with oscillations in velocity and line width.  Further active region velocity studies with SUMER were performed by \\citet{marsch04} and \\citet{doschek06}. The former presented Doppler maps for three active regions \\citep[including that of][]{winebarger02} and, although not specifically remarked on, it is clear that several fan loop structures (including some close to a sunspot) in the data display significant redshifts in \\ion{Ne}{viii} \\lam770. The loops also appear to show redshifts in \\ion{C}{iv} \\lam1548. \\citet{doschek06} presented Doppler maps in \\ion{C}{iv} \\lam1548, \\ion{S}{v} \\lam786 and \\ion{Ne}{viii} \\lam770 (formed at 0.1, 0.2 and 0.7~MK, respectively) for an active region and two ``plumelike'' structures (almost certainly the same as fan loops) were shown to exhibit redshifts of 5--10~\\kms\\ in \\ion{S}{v} and \\ion{Ne}{viii} although in this case not in \\ion{C}{iv}. The SUMER instrument had limited access to coronal emission lines, and no velocity studies of fan loops were made for higher temperatures than \\ion{Ne}{viii} (0.7~MK).   The EUV Imaging Spectrometer \\citep[EIS;][]{culhane07} on board the \\hinode\\ satellite \\citep{kosugi07}  is the first solar ultraviolet spectrometer to routinely allow the measurement of coronal emission line Doppler shifts to precisions of 0.5~\\kms\\ \\citep{mariska08}. It also has access to lines from the upper transition region (temperatures 0.2--1.0~MK) and thus there is some temperature overlap with the SOHO/SUMER instrument. EIS therefore, for the first time, allows the change in velocity structure from the transition region to the corona to be investigated with a single instrument. Studies of \\citet{delzanna08} and \\citet{tripathi09} have demonstrated that specific locations in an active region can show a change from redshifts to blueshifts from the transition region to the corona. \\citet{delzanna09a,delzanna09b} presented velocity measurements in fan loops near a sunspot and found evidence of decreasing redshift with temperature: \\ion{Fe}{vii} at $+30$~\\kms, \\ion{Fe}{viii} at $+20$~\\kms\\ and \\ion{Fe}{ix} at $+10$~\\kms\\ (lines formed at 0.6, 0.7 and 0.8~MK, respectively). More recently, \\citet{warren11} have demonstrated that the velocity structure within the vicinity of active region fan loops changes with temperature. They presented velocity maps in different emission lines showing that the fan loops display redshifts (downflows) for temperatures up to 0.8~MK. There is an abrupt change to approximately zero velocity at a temperature of 1.0~MK, and then a change to strong blueshifts at higher temperatures of 1.0--2.0~MK. These blueshifts represent the active region outflow regions that have been described in earlier EIS papers \\citep{delzanna08,harra08,doschek08}. The authors suggested that the cool fan loops are distinct structures from the active region outflows, with the latter probably representing open field plasma that escapes to the solar wind, while the former are closed loops.  Further support for a distinction between the outflow regions and the fan loops is provided by \\citet{ugarte11} who compared time sequences of images and velocity maps from the EIS lines \\ion{Si}{vii} \\lam275.37 ($T=0.7$~MK) and \\ion{Fe}{xii} \\lam195.12 (1.6~MK).  Inspection of the  images presented in \\citet{warren11} shows that the fan loops have a filamentary structure in both velocity and intensity and thus there is the possibility that the bright (in intensity) filaments within the loop may not necessarily correspond to filaments with the strongest velocity shifts. The present work considers a specific example of a fan loop that is distinct in intensity and that is aligned along the line-of-sight to the observer (thus making the measured line-of-sight velocities a good approximation to the actual loop velocities). The aim is to place the qualitative observations of \\citet{warren11} onto a quantitative basis. A key part of the work is to make a general prescription for deriving absolute velocities with correct error bars with EIS, which has not been previously done. ",
        "conclusions": "\\label{sect.discuss}  The most significant result from this work is that a single loop structure observed with EIS exhibits significantly different line-of-sight velocities in two emission lines that are relatively close in temperature. \\ion{Fe}{viii} \\lam185.21 shows a redshift of $\\approx 15$~\\kms\\, while \\ion{Fe}{x} \\lam184.54 shows velocities close to 0~\\kms. The differences apply at the same spatial locations in the loop, and are sustained over a projected distance of 9~Mm along the length of the loop. An immediate consequence is that the loop consists of at least two plasma components at different temperatures: one is cool and downflowing, the other is hot and at rest. The fact that the velocity difference is maintained over a large spatial distance means that the two components represent two distinct structures rather than, e.g., the loop simply showing downflows near the bottom, cooler part of the loop and stationary plasma in the upper, hotter part of the loop. The higher spatial resolution of the TRACE instrument revealed that the EIS loop probably consists of at least two narrow `strands' visible at 1~MK in the 171 filter (Fig.~\\ref{fig.eis-ims}). However, since both strands are emitting at 1~MK then they will both contribute to the EIS \\ion{Fe}{x} \\lam184.54 emission line and so do not represent the two velocity components.   This result has important consequences if it is treated as a general property of fan loops. Consider, for example, the work of \\citet{winebarger02} who found \\ion{Ne}{viii} redshifts in fan loops identified from TRACE 171 images. The authors proceeded to make a loop model with a steady flow that could reproduce the TRACE 171 intensity profile along the loop. Since \\ion{Ne}{viii} is formed at the same temperature as \\ion{Fe}{viii} then the redshift result is consistent with the present work \\citep[although the magnitude of the velocity was significantly larger in the][loop system]{winebarger02}. The TRACE 171 channel is dominated by \\ion{Fe}{x} with a smaller contribution from \\ion{Fe}{ix}. If our EIS result is extrapolated to the \\citet{winebarger02} loop system, then the TRACE 171 emission has \\emph{little or no physical relation} to the \\ion{Ne}{viii} emission, and so the \\ion{Ne}{viii} velocity result can not be included in a model that seeks to interpret the intensity variation in the 171 loop image.  One concept for how the loop may be physically structured is shown in Fig.~\\ref{fig.cartoon}. Two distinct loop types (or strands) are shown. Type 1  is stationary and emits \\ion{Fe}{x} along a large portion of its length, while type 2 is downflowing and emits \\ion{Fe}{viii} along a large portion of its length. The ``centroid'' of the \\ion{Fe}{x} emission in the type 1 strand is located somewhat higher than that of  \\ion{Fe}{viii} in the type 2 strand, as suggested by the EIS images (Fig.~\\ref{fig.eis-ims}). Since the type 1 strand must cool before descending into the photosphere, it must also emit \\ion{Fe}{viii} emission, but this emission will be much more compact than that coming from the type 2 strand. We note that the velocity of \\ion{Fe}{viii} \\lam185.21 does fall at the lowest heights in the EIS loop (Fig.~\\ref{fig.loopb-vel}) and the width increases (Fig.~\\ref{fig.loopb-wid}) which may indicate that the type 1 strand is contributing significant \\ion{Fe}{viii} emission.  One question relates to the presence of any emission in type 2 strands from ions hotter than \\ion{Fe}{viii}. It would be surprising if the strand did not achieve coronal temperatures, yet the coronal emission lines observed by EIS show that such emission either does not partake in the \\ion{Fe}{viii} flow, or that it is sufficiently weak that it is masked by the stationary coronal plasma of the type 1 strands.  \\begin{figure}[h] \\epsscale{0.7} \\plotone{f14.eps} \\caption{A conceptualization of the two types of loop strand that could explain the EIS velocity results. See text for more details.} \\label{fig.cartoon} \\end{figure}  The work of \\citet{delzanna08}, \\citet{tripathi09}, \\citet{delzanna09a} and \\citet{warren11} suggests that the velocity results found here may be a common feature of fan loops. If this is correct then, despite the division of the loop into two types of independent strand, these two strands always appear together in fan loops, and their relative distribution/strength always conspires to make \\ion{Fe}{x} appear stationary and \\ion{Fe}{viii} downflowing. If there was significant imbalance towards type 1 strands, then \\ion{Fe}{viii} would no longer display extended emission, instead being compressed into a small spatial region at the base of the \\ion{Fe}{x} strand. While an imbalance towards type 2 strands would yield a fan loop with either little \\ion{Fe}{x} emission, or with a \\ion{Fe}{x} line profile showing significant downflow velocity. A natural solution to the apparent connectedness of the two strand types is that the type 2 strand represents a later, cooling stage in the evolution of a type 1 strand. However, the number of stationary strands is well-balanced with the number of cooling, downflowing strands in order that a predominance of one over the other is not seen in the observations.  These arguments are somewhat speculative and require a survey of many loops before they can be confirmed. Certainly a check on the velocity difference between \\ion{Fe}{viii} and \\ion{Fe}{x} needs to be performed on many loops from different active regions. A thorough velocity analysis such as the one presented in this work is not necessary. Simply measuring the wavelength separation of the nearby \\ion{Fe}{viii} \\lam185.21 and \\ion{Fe}{x} \\lam184.54 lines in the loops and comparing with the standard, off-limb rest separation will be sufficient. A study in the same loops of the intensity distributions of the \\ion{Fe}{viii} and \\ion{Fe}{x} lines with height would also be profitable by determining the heights of maximum emission. A time history of how the \\ion{Fe}{viii}--\\ion{Fe}{x} velocity difference changes (or does not change) for a loop is also an important study, as it may reveal if the balance between type 1 and type 2 strands changes with time.   The concept of apparently monolithic loop structures consisting of multiple, unresolved strands has been explored previously by theorists, particularly with regard the nanoflare theory of coronal heating \\citep{parker88,cargill94}, where small heating events only heat individual strands, but combined can account for global loop properties such as temperature, density and emission measure. Multi-stranded loop models have also been cited to explain observations that appear to show some loops have an isothermal temperature profile \\citep{lenz99,reale00}.  Dynamic, 1D models of individual loops/strands that include plasma flows have been made by \\citet{spadaro03}, \\citet{spadaro06} and \\citet{bradshaw10}. \\citet{spadaro03} considered active region loops of 40--80~Mm in length, with peak temperatures of around 4~MK and coronal densities of up to $2\\times 10^{10}$~cm$^{-3}$; while \\citet{spadaro06} modeled quiet Sun loops with lengths of 5--10~Mm, peak temperatures up to 1.3~MK, and coronal densities up to $10^9$~cm$^{-3}$. The parameters of the loop modeled by \\citet{bradshaw10} are closest to  those of the loop studied here -- a loop length of 67~Mm, peak temperature of 3.6~MK and density of $3\\times 10^9$~cm$^{-3}$ -- so we summarize briefly the velocity results from this work. The loop was allowed to reach a hydrostatic equilibrium with a coronal apex temperature and cool footpoints, and then allowed to cool with no heating applied. Enthalpy and conduction fluxes balance the radiative losses of the plasma, and \\citet{bradshaw10} found that the enthalpy flux dominates in certain temperature regimes and evolution periods. As the loop progressively cools, increasing downflows are found near the peak temperature of the loop, varying from $\\approx$10~\\kms\\ around 2~MK to $\\approx$40~\\kms\\ at 0.3~MK. At the temperatures at which \\ion{Fe}{viii} and \\ion{Fe}{x} are formed (0.7--1.1~MK) the velocities are around 15--20~\\kms, and most of the length of the loop shows significant downflows. The \\ion{Fe}{viii} velocity measured in ``loop structure B'' of the present work is consistent with this model, but the flow is only found in the lower portion of the loop, and \\ion{Fe}{x} shows no significant velocity.  The model can not be directly compared with the observation as the observed loop likely consists of many strands, and so a better comparison would be with a time-average of cooling strands randomly distributed in time. Even then, there is likely a contribution to the emission line intensities and velocities from the heating phase of the strands which is not modeled by \\citet{bradshaw10}. The \\ion{Fe}{viii}--\\ion{Fe}{x} velocity result presented here would seem to be a valuable constraint on loop models of this sort, and modelers are encouraged to investigate how it can be reproduced.     Finally, a comment on coronal velocity measurements. The present paper has laid out the difficulties of determining absolute velocities from the EIS instrument. Although EIS is capable of measuring velocities as \\emph{precise} as 0.5~\\kms\\ \\citep{mariska08}, they are only \\emph{accurate} to, at best, 4~\\kms. Improved accuracy at EUV wavelengths (i.e., below 912~\\AA) is unlikely for the forseeable future of solar physics due to the lack of absolute wavelength standards. The only solution appears to be observing coronal lines above the Lyman limit where either a calibration lamp can be used, or photospheric/low chromosphere lines with small velocities that serve as wavelength fiducials can be observed. Since the allowed transitions of all coronal ions are found below 700~\\AA\\ the emission lines must be observed in second, third or even fourth spectral order, but then blending can become a significant issue. There are some coronal forbidden lines that are found above the Lyman limit but below where the photospheric continuum begins to become significant ($\\approx 1400$~\\AA). The most notable such line is probably \\ion{Fe}{xii} \\lam1242."
    },
    "1107/1107.0367_arXiv.txt": {
        "abstract": "We analyze the generation of polarization cross-talk in Stokes polarimeters by atmospheric seeing, and its effects on the noise statistics of spectro-polarimetric measurements for both single-beam and dual-beam instruments. We investigate the time evolution of seeing-induced correlations between different states of one modulation cycle, and compare the response to these correlations of two popular polarization modulation schemes in a dual-beam system. Extension of the formalism to encompass an arbitrary number of modulation cycles enables us to compare our results with earlier work. Even though we discuss examples pertinent to solar physics, the general treatment of the subject and its fundamental results might be useful to a wider community. ",
        "introduction": "Many important solar phenomena are driven by the interaction of turbulent plasma with magnetic fields. Measurements of the full state of polarization of solar spectral lines are routinely used to infer properties of solar vector magnetic fields from their imprints in the emergent spectra through the Zeeman and Hanle effects. All ground-based measurements are detrimentally affected by image distortion introduced by atmospheric seeing, which has significant power at frequencies up to at least $100\\,$Hz, significantly higher than the frame rates of typical CCD cameras. The seeing-induced errors due to image distortion during an observation are expected to increase as the telescope aperture increases in size, since the telescope can then resolve finer structures and hence observe steeper gradients in images.  The Advanced Technology Solar Telescope \\cite[ATST;][]{Ri08}, a 4-m aperture, off-axis telescope, will soon enter construction phase. The European Solar Telescope \\cite[EST;][]{Co08} is also a 4-m class telescope currently under design. Spectro-polarimetry is the prime mode of operation of this new generation of solar telescopes, but there is an urgent need to identify potential pitfalls before committing to a particular design for the modulation and detection of polarized light using such telescopes. Space-based spectro-polarimetric instruments, like those on-board the Hinode \\cite[Solar-B;][]{Ts08} and the Solar Dynamics Observatory \\cite[SDO;][]{Sc12} spacecrafts, are by necessity fed by smaller telescopes, and while their observations profit greatly from the absence of atmospheric disturbances, they are still affected by residual image motion due to spacecraft jitter.  The purpose of the present paper is to re-examine and extend earlier theoretical studies of polarization errors introduced by the effects of atmospheric seeing. Potentially, this can help refine the instrument requirements for the polarization systems of these large, multi-national projects, as well as of future solar space missions, such as Solar-C \\citep{Sh11}.  Two modes of operation are of particular interest for our study. Slit-based spectro-polarimeters often use long integration times, including a number of modulation cycles that is typically of the order of a few tens. Because of the need to scan across the spectral line domain, tunable imaging spectro-polarimeters typically implement much fewer measurements of the modulated intensities -- often just one modulation cycle -- for each wavelength position. The formalism developed in this paper is general, and can be applied to both types of instruments and observations. A substantial difference between the two types of observations is that post-processing techniques such as image shifting and de-stretching can be effectively adopted to partially correct for seeing-induced distortion in observations done with imaging polarimeters, whereas this option is not available in the case of slit-based spectro-polarimeters.  A separate question -- also of fundamental importance for observational spectro-polarimetery -- concerns the effects of polarization calibration errors on the degree of polarization cross-talk affecting an observation. This subject has been exhaustively treated by \\cite{AC08}, and therefore it will not be addressed here.  The plan of this paper is the following. In Sect.~\\ref{sec:seeing}, we describe a model for the effects of atmospheric seeing on the polarization signals detected by an observer. In Sect.~\\ref{sec:formalism}, we lay the foundation of the formalism. Using the point of view of the statistics of random processes we define the fundamental statistical observables for spectro-polarimetry in the presence of atmospheric seeing. In Sect.~\\ref{sec:dualbeam}, this formalism is extended to the case of dual-beam polarimetry, emphasizing those aspects of the problem  where this extension is not trivial. In Sect.~\\ref{sec:covariances} we make use of the conceptual framework of the statistics of stationary random processes to gain a deeper understanding of the effects of seeing correlations on spectro-polarimetric measurements. In particular, in that section we study qualitatively the effect of adaptive optics (AO) corrections on the temporal decay of seeing correlations. Those results are successively applied to study the performance of two popular modulation schemes in the presence of seeing. In Sect.~\\ref{sec:generalization}, these results are further generalized to the case in which seeing correlations extend over the duration of one elemental observation (e.g., one slit position), and the effects of seeing on the performance of optimally efficient modulation schemes are studied both in the absence and in the presence of low-order (i.e., tip-tilt) AO corrections (based on observed data). Finally, we compare our results and conclusions with those from previous works on the subject. ",
        "conclusions": "\\label{sec:earlier}  In this paper, we have approached the determination of seeing-induced cross-talk noise from a statistical point of view. Equations (\\ref{eq:av_signal_i})--(\\ref{eq:sigma_modint}) and (\\ref{eq:av_S_i}) and (\\ref{eq:var_S_i}) form the basis of our derivation (see Sect.~\\ref{sec:formalism}), so it is important to compare those results with previous work \\citep{Li87,Ju04}.  In the work of \\cite{Li87} and \\cite{Ju04}, the ``variances'' there defined depend explicitly on the total observation time and the modulation frequency, through the integral of the product of the seeing power spectrum with a real function of frequency that depends implicitly on both (cf.~Eq.~[15] of \\citealt{Li87}, and points 1--5 of Sect.~2 of \\citealt{Ju04}). In our formalism, the total integration time does not appear explicitly because Eq.~(\\ref{eq:sigma_modint}) represents the variance on a \\textit{single} measurement of $\\mathscr{I}_i$. However, if a number $N$ of such measurements are made, \\textit{and those measurements can be considered statistically uncorrelated}, then the variance on the \\textit{average} signal $\\bar{\\mathscr{I}}_i$ is reduced by a factor $N$. In such case, for a fixed modulation frequency, but increasing the total observation time, we expect the same qualitative scaling of variance with the integration time as derived in previous work. Secondly, the dependence on the modulation frequency and the seeing power spectrum is also apparent within our formalism -- already in the case of a single measurement -- when we consider Eq.~(\\ref{eq:signal_i}).  Incidentally, we note how the power spectrum of the seeing appears most naturally in our approach as the Fourier transform of the two-time correlation function of the seeing displacement vector (see Sect.~\\ref{sec:covariances}), whereas in previous work it is identified instead with the modulus square of the Fourier transform of the seeing displacement.\\footnote{We observe that the ordinary Fourier transform is not well defined for the seeing displacement vector, because the associated random process is not limited in time. One can get around this problem by defining a generalized Fourier transform of the seeing, as the limit of ordinary Fourier transforms of finite samples of the seeing for increasing duration of those samples \\citep[e.g.,][]{MW95}.} The correspondence between these two approaches to the definition of the power spectrum of a stationary random process is clearly described by \\cite{MW95}.  The importance of the covariances $E([\\mathscr{I}_j-\\bar{\\mathscr{I}}_j][\\mathscr{I}_k-\\bar{\\mathscr{I}}_k])$ in typical cases has fundamental implications for the concept of polarimetric measurements. Through our analysis we are able to quantify the significance of seeing-induced correlations between measurements corresponding to different modulation steps, and how these correlations decay for decreasing modulation frequencies. Based on those results, it must be expected that seeing-induced correlations between different modulation states typically extend beyond the time interval of one modulation cycle. In other words, the Stokes vector measurements corresponding to different modulation cycles during an elemental observation are in general statistically correlated. Under this condition, we must expect that the variance of the average signal $\\bar{\\mathscr{I}}_i$ will obey the ordinary scaling law by the total number $N$ of cycles of the elemental observation only approximately. Formally, one should consider instead such elemental observation as a \\textit{single} measurement that is realized through the totality of the $N$ modulation cycles, and thus characterized by a corresponding $nN\\times 4$ modulation matrix. The variance of this measurement can then be determined through the usual equations, adopting such extended definition of the modulation and demodulation matrices (see Sect.~\\ref{sec:generalization}). A similar approach is taken also in the earlier work \\citep{Li87,Ju04}, where these correlations are made manifest by performing a spectral analysis of the demodulated signal over the entire time interval of the elemental observation.  Our formalism reveals however that, in the work of \\cite{Li87} and \\cite{Ju04}, the covariance terms corresponding to contributions proportional to $\\bm{\\nabla}S_i\\cdot\\bm{\\nabla}S_j$, with $i\\ne j$ (see Eq.~[\\ref{eq:first1}]), are not accounted for. In fact, \\cite{Li87} derives the variances as Eq.~(12) of that paper directly from Eq.~(11), which are explicitly of diagonal form. In other words, looking at our Eqs.~(\\ref{eq:first2}) and (\\ref{eq:var_S_matform}), previous work has computed the seeing-induced noise always under the assumption that the gradient matrix $\\tens{G}$ was diagonal. The additional off-diagonal components would arise in the development by \\cite{Li87} and \\cite{Ju04} when properly evaluating the expectation value of $(O_r - \\bar O_r)^2$ in the notation of \\cite{Li87}. Equations (\\ref{eq:first2}) and (\\ref{eq:var_S_matform}) also clarify the vanishing of the cross-talk terms in the limit of large modulation frequencies ($\\chi\\Delta t\\to 0$). Because all the integrals $\\mathscr{T}_{|j-k|}(\\Delta t)$ tend to the same value in such limit, the vector of the Stokes variances becomes simply proportional to the diagonal of the gradient matrix, $\\tens{G}$."
    },
    "1107/1107.0017_arXiv.txt": {
        "abstract": "We make detailed theoretical predictions for the  assembly properties of the Local Group (LG) in the standard $\\Lambda$CDM cosmological model. We use three cosmological N-body dark matter simulations from the Constrained Local Universe Simulations (CLUES) project, which are designed to reproduce the main dynamical features  of the matter distribution down to the scale of a few Mpc around the LG. Additionally, we use the results of an unconstrained simulation with a  sixty times larger volume to calibrate the influence of cosmic variance.  We  characterize the Mass Aggregation History (MAH) for each halo by three characteristic times, the formation, assembly and last major merger times.  A major merger is defined by a minimal mass ratio of $10:1$.  We find that the three LGs share a similar MAH with formation and last major merger epochs placed on average $\\approx 10-12$ Gyr ago. Between $12\\%$ and $17\\%$ of the halos in the mass range $5\\times 10^{11}\\hMsun <M_{h} < 5\\times 10^{12}\\hMsun$ have a similar MAH. In a set of pairs of halos within the same mass range, a fraction of $1\\%$ to $3\\%$ share similar formation properties as both halos in the simulated LG. An unsolved question posed by our results is the dynamical origin of the MAH of the LGs.  The isolation criteria commonly used to define  LG-like halos in unconstrained simulations do not narrow down the halo population into a set with quiet MAHs, nor does a further constraint to reside in a low density environment.  The quiet MAH of the LGs provides a favorable environment for the formation of disk galaxies like the Milky Way and M31. The timing for the beginning of the last major merger in the Milky Way dark matter halo matches with the gas rich merger origin for the thick component in the galactic disk.  Our results support the view that the specific large and mid scale environment around the Local Group play a critical role in shaping its MAH and hence its baryonic structure at  present. ",
        "introduction": "Observations of the Milky Way (MW) and the galaxy M31 shape to a great extent our understanding of galaxy formation and evolution. In particular, three landmarks have been pivotal in the development of theoretical studies of structure formation: a) the abundance of MW galaxy satellites that motivated one of the strongest points of tension with the now-standard $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) paradigm of structure formation \\citep{1999ApJ...522...82K,1999ApJ...524L..19M}, b) the spatial distribution of the same satellites which triggered discussions on how unique the host dark matter halo of the MW is \\citep{2009MNRAS.394.2223M} and c) the measurements of the tidal debris of disrupted merging galaxies around the MW and M31 galaxy, confirming the  hierarchical nature of galaxy evolution, one of the fundamental characteristics of $\\Lambda$CDM \\citep{2009Natur.461...66M}. However, inferring general conclusions on galaxy evolution based on observations of these two galaxies requires an assessment on how biased the properties of the MW and M31 are with respect to a given control population.   In the framework of $\\Lambda$CDM, the study of the MW and M31 starts by modeling their individual host dark matter halos, \\emph{assuming} that their simulated formation histories are \"typical\", or at least compatible with the assembly of the real Local Group \\citep{2009MNRAS.395..210D,2010MNRAS.406..896B}.  The basic definition of a LG (in terms of the dark matter distribution) has two basic elements based on the state of the system \\emph{today}: (i)  the estimated masses of the dark matter halos corresponding to the MW and M31  (see for instance \\cite{2010MNRAS.406..264W} and references therein) and (ii) the isolation of these two halos from other massive structures \\citep{2004AJ....127.2031K}. Two additional constraints could be the separation and the relative velocity of the two halos \\citep{2005ApJ...635L..37R}. However, the condition on the LG isolation admits a strict formulation, by requiring that the environment, in terms of the mass and position of the dominant galaxy clusters in the Local Universe, be as close as possible to the one inferred from observations. Such an additional condition imposes restrictions on the possible outcomes of structure formation on scales of the order of $\\sim5$ Mpc. This is considered here as the meso-scale as opposed to the large ($\\gtrsim 5$ Mpc) or the small ($\\lesssim 1$ Mpc) scales.    The new feature in the analysis presented in this paper is the inclusion of such observational constraints around the LG environment in the initial conditions of the simulation. In a series of three simulations from such initial conditions, in a WMAP5 cosmology with a normalization $\\sigma_{8}=0.817$ \\citep{2009ApJS..180..330K}, we are able to define a sample of three LG dark matter halo pairs that form and evolve under specific conditions reflecting  structure of the Local Universe. In addition we will take advantage of one of the largest cosmological simulations carried out to date, the Bolshoi Simulation \\citep{2010arXiv1002.3660K}, to explore a larger sample of halos within the mass range of the LG, and calibrate possible cosmic variance effects.   We analyze the constrained simulations with the primary goal of quantifying the assembly histories of the LG halos. This is driven by two different motivations. One is to find out whether the simulated LGs, that are selected by dynamical considerations pertaining to their redshift zero structure, have mass aggregation histories (MAHs) that lead to the formation of disk galaxies like the MW and M31.  The other is to find out whether such a MAH is dictated by by meso-scale environment of the LG, or whether a random  selection of objects similar to the LG is likely to have a similar MAH.   In Section \\ref{sec:clues}, we describe our simulations and the method to re-construct the mass aggregation histories. In Section \\ref{sec:sample} we describe how we build the different control samples for our statistical analysis. In Section \\ref{sec:results} we study the MAHs in the different samples and argue that the selection by different isolation criteria does not induce a strong bias in the statistics describing the MAHs. In Section \\ref{sec:discussion} we discuss the possible origin of these findings and comment on the connection with observations of the MW and M31. In Section \\ref{sec:conclusions} we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We use constrained simulations of the local Universe to study the dark matter mass aggregation history (MAH) of the Local Group (LG). Two basic questions motivate this study: 1. To what extent the simulated LGs can account for the observed structure of the MW and M31 galaxies? Namely, if the disk dominated morphology implies that the MW and M31 halos had a quiet MAH over the last $\\approx 11$Gyrs, can simulations recover this recent quiet history? 2. Does this quiet MAH arise from the intrinsic properties of the DM halos, or is it induced by environment within which the LG is embedded? Is the implied MAH of the LG triggered by the large and meso-scales, or is it induced by the small, i.e. galactic and sub-galactic, scales?   The methodology adopted here is to use constrained simulations of the local Universe, designed to reproduce the large and meso-scales of the LG environment, and search for halos that resemble the actual LG. The identification of a pair of halos as a LG-like object is based on a set of isolation and dynamical criteria, all formulated by their redshift zero structure, in complete ignorance of their formation history. A LG-like object that is found close to the actual position of the observed LG with respect to the large scale structure environment is defined here as a LG. By construction a constrained simulation can have only one LG or none at all. Indeed, out of a suit of 200  constrained simulations only 3 harbour a LG. Controlled samples of individual halos and pairs have been constructed as reference samples. The analysis has been extended to the unconstrained Bolshoi simulation that is used here for an unbiased reference \\citep{2010arXiv1002.3660K}.    The construction of the identification of the 3 LGs is done independently of the MAH of the halos. Yet, the MW's and M31's halos of the 3 LGs all have a common quiet MAH, defined as having the last major merger, formation and assembly look-back time extending over $\\approx (10\\ -12)$ Gyr. This  quiet formation history of the simulated LGs can help to explain the disk dominated morphology of the MW and M31, adding evidence to the internal instability origin of the spheroidal component of the MW \\citep{2010ApJ...720L..72S}.  Based on measurements of the eccentricity of orbits in the MW, it has been recently claimed \\citep{2009MNRAS.400L..61S,2010ApJ...725L.186D,2011MNRAS.tmp..260W}, that a  rich merger taking place 10.5 to 8 Gyr ago is a favoured mechanism  explain the thick disk in the MW  \\cite{2004ApJ...612..894B}.  Our finding of a quiet MAH of the LG provides a suitable platform for such a process to take place.    The LG halos are assumed here to be selected from FOF halos in the mass range $5\\times 10^{11}\\hMsun <M_{h} < 5\\times 10^{12}\\hMsun$ at $z=0$. Between $12\\%$ and $17\\%$ of these halos are found to have a quiet MAH, depending on the detailed definition of the quiet parameter space. From this point of view the MW and M31 halos are not  rare. However, how likely is a pair of halos to have such a quiet history, shared by both halos? Making the naive null assumption that the MAH of a halo is an intrinsic property of a halo independent of its environment then the fraction of pairs should be the product of the fractions for a single halo. Indeed, the {\\it Pairs} sample drawn out of the Bolshoi simulation confirms this assertion, finding that between $1\\%$ to $3\\%$ of the pairs have as quiet an MAH as the LG systems do. The probability of having selecting 3 pairs randomly and finding them with a quiet MAH is on the order of $\\sim 10^{-5}$.  Next, we look for what dynamical or environmental property determines rethe MAH of a LG-like object.  We find here  that the mere pairing of the MW-like halos does not affect the MAH fiducial times. Imposing the isolation and dynamical constraints that define the {\\it Isolated Pairs} sample  does not affect it either. This leaves us with an open question as to what determines the MAH of halo pairs similar to the LG. The one hint that we have is that all of the three LGs reside in the same large and meso-scale environment. We speculate that the cosmic web plays a major role in shaping the MAH of LG-like objects,  although it is not yet clear what mechanism is responsible. A larger sample of constrained LGs is needed to confirm and further explore the reasons behind this result."
    },
    "1107/1107.0826_arXiv.txt": {
        "abstract": "We present a numerical extension to the analytical propagation model introduced in \\citet{hein08} to describe the leptonic population in the galactic disc. The model is used to derive a possible identification of the components that contribute to the leptonic cosmic ray spectrum, as measured by PAMELA, Fermi and HESS, with an emphasis on secondary $e^+-e^-$ production in collisions of cosmic ray particles with ambient interstellar medium (ISM). We find that besides secondaries, an additional source symmetric in $e^+$ and $e^-$ production is needed to explain both the PAMELA anomaly and the Fermi bump, assuming a power-law primary electron spectrum. Our model also allows us to derive constraints for some properties of the ISM. ",
        "introduction": "Recently the leptonic component of the cosmic ray spectrum has gained new attention. New observations from ATIC, PAMELA and Fermi show a deviation from a power-law in the form of an excess in both the electron and positron spectra. Annihilating dark matter \\citep{allahverdi08} and nearby pulsars \\citep{busching08,grimani2009,amatoblasi2010} (among other hypotheses) have been proposed as possible sources of the excess leptons. Regardless of the source, a propagation model is needed to connect the energy spectrum measured on Earth with the injection \\mbox{spectra.}  We present our numerical cosmic ray transport model in application to the high energy electron transport in the ISM. Spatial and momentum diffusion, particle escape, acceleration via Fermi I and continuous energy losses were taken into account and their effects on the steady-state energy spectrum analyzed. In solving the transport equation we employed quasi-linear transport theory, the diffusion approximation and a separation of the spatial and momentum problem to obtain the leaky-box-equation, which was then solved numerically. The spatial problem was solved analytically in cylindrical and prolate spheroidal coordinates.  The transport model was employed to calculate the spectrum of secondary electrons in our galaxy. We assume the leptons from pioArXivn decay to be the dominating component of the secondary spectrum. The lepton spectrum from pion collisions of highly relativistic cosmic ray protons with thermal protons in the ISM was calculated using the parametrization of pion production in p-p-collisions from \\citet{Aharonian} and used as the injection spectrum for the transport model. With realistic simulation parameters the resulting positron flux in the vicinity of the solar system lies remarkably close to the low-energy PAMELA data. Assuming a generic power-law primary $e^-$ spectrum and an additional $e^+-e^-$ source allows us to fit the datasets from PAMELA, Fermi and HESS very nicely and to put constraints on several transport model parameters and the corresponding ISM properties. ",
        "conclusions": ""
    },
    "1107/1107.2538_arXiv.txt": {
        "abstract": "We construct a domain wall solution in $F(R)$ gravity. We reconstruct a static domain wall solution in a scalar field theory. We also reconstruct an explicit $F(R)$ gravity model in which a static domain wall solution can be realized. Moreover, we show that there could exist an effective (gravitational) domain wall in the framework of $F(R)$ gravity. In addition, it is demonstrated that a logarithmic non-minimal gravitational coupling of the electromagnetic theory in $F(R)$ gravity may produce time-variation of the fine structure constant which may increase with decrease of the curvature, and that this model would be ruled out by the constraints on the time variation of the fine structure constant from quasar absorption lines. We also present cosmological consequences of the coupling of the electromagnetic field to a scalar field as well as the scalar curvature and discuss the relation between variation of the fine structure constant and the breaking of the conformal invariance of the electromagnetic field. ",
        "introduction": "According to recent cosmological observations, e.g., Supernovae Ia (SNe Ia)~\\cite{SN1}, cosmic microwave background (CMB) radiation~\\cite{WMAP, Komatsu:2010fb}, large scale structure (LSS)~\\cite{LSS}, baryon acoustic oscillations (BAO)~\\cite{Eisenstein:2005su}, and weak lensing~\\cite{Jain:2003tba}, it has been implied that the current expansion of the universe is accelerating. Studies on the late time cosmic acceleration are classified into the representative two categories. One is to introduce dark energy such as cosmological constant in the framework of general relativity (for a recent review, see, e.g.,~\\cite{Li:2011sd}). The other is to modify the gravitational theory, for instance, $F(R)$ gravity, where $F(R)$ is an arbitrary function of the scalar curvature $R$ (for recent reviews on $F(R)$ gravity, see, e.g.,~\\cite{Review-Nojiri-Odintsov, Book-Capozziello-Faraoni, Clifton:2011jh}).  Recently, not only temporal~\\cite{Time-variation, Murphy:2003hw} but also spatial~\\cite{Webb:2010hc} variations of the fine structure constant ${\\alpha}_{\\mathrm{EM}}$ have been suggested. To account for the spatial variation of ${\\alpha}_{\\mathrm{EM}}$, the signature of a domain wall produced in the spontaneous symmetry breaking involving a dilaton-like scalar field coupled to electromagnetism has been considered in Ref.~\\cite{Olive:2010vh}. Furthermore, in Ref.~\\cite{Chiba:2011en} it has been shown that a runaway domain wall, which is formed by a runaway type potential of a scalar field~\\cite{Cho:1998jk}, can explain both the time variation by its potential and the spatial one by its formation simultaneously. It is interesting to note that in Ref.~\\cite{M-B-C} time and spatial variations of coupling constant have been studied and that when the chameleon field is introduced, variations of coupling constant is related to the chameleon mechanism~\\cite{Chameleon-mechanism}.   On the other hand, a domain wall solution in the framework of $F(R)$ gravity has not been investigated in detail yet. In particular, it is interesting to reconstruct an $F(R)$ gravity model in which a domain wall solution can be realized. It is known that $F(R)$ gravity can be written as a corresponding scalar field theory through a conformal transformation to the Einstein frame. In this paper, we reconstruct an explicit $F(R)$ gravity model in which a static domain wall solution can be realized. First, by using a procedure proposed in Ref.~\\cite{Capozziello:2005tf}, we reconstruct a static domain wall solution in a scalar field theory. Next, in a similar configuration, we reconstruct an explicit form of $F(R)$ with forming a static domain wall solution. Moreover, by applying the method of reconstruction of $F(R)$ gravity in Ref.~\\cite{Reconstruction-F(R)-N-O}, we show that there could exist an effective (gravitational) domain wall in the framework of $F(R)$ gravity. In addition, we discuss an issue of a connection between $F(R)$ gravity and variation of the fine structure constant by exploring non-minimal Maxwell-$F(R)$ gravity. Furthermore, we present cosmological consequences of the coupling of the electromagnetic field to a scalar field as well as the scalar curvature. We also study the relation between variation of the fine structure constant and the breaking of the conformal invariance of the electromagnetic field. 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. Moreover, in terms of electromagnetism we adopt Heaviside-Lorentz units.  The paper is organized as follows. In Sec.\\ II, we describe $F(R)$ gravity and a corresponding scalar field theory by using a conformal transformation of $F(R)$ gravity to the Einstein frame. In Sec.\\ III, we reconstruct a static domain wall solution in a scalar field theory. In Sec.\\ IV, we also reconstruct an explicit $F(R)$ gravity model in which a static domain wall solution can be realized. In Sec.\\ V, we demonstrate that there could exist an effective (gravitational) domain wall in $F(R)$ gravity. In Sec.\\ VI, we consider non-minimal Maxwell-$F(R)$ gravity and examine a relation between $F(R)$ gravity and variation of the fine structure constant. In addition, we investigate cosmological consequences of the coupling of the electromagnetic field to a scalar field as well as the scalar curvature in Sec.\\ VII. Finally, conclusions are given in Sec.\\ VIII. ",
        "conclusions": "In the present paper, we have studied a domain wall solution in $F(R)$ gravity. We have reconstructed a static domain wall solution in a scalar field theory. We have also reconstructed an explicit $F(R)$ gravity model in which a static domain wall solution can be realized. Furthermore, we have shown that there could exists an effective (gravitational) domain wall in the framework of $F(R)$ gravity. Moreover, it has been illustrated that a logarithmic non-minimal gravitational coupling of the electromagnetic theory in $F(R)$ gravity may produce time-variation of the fine structure constant which may increase as the curvature decreases. In addition, we have described cosmological consequences of the coupling of the electromagnetic field to not only a scalar field but also the scalar curvature and remarked the relation between variation of the fine structure constant and the breaking of the conformal invariance of the electromagnetic field.  The reconstruction technique was applied here to inducing of domain wall solution in modified gravity (cf. the case of black hole reconstruction in Ref.~\\cite{Nojiri:2006jy}). It is clear that similar methods may be applied to generation of other solutions in modified gravity, like topological defects, cosmic strings, etc. This question will be discussed elsewhere."
    },
    "1107/1107.2012_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\label{sect:final}   Although it is reasonably feasible to compare the major properties of a Milky Way galaxy formed in a $\\Lambda$CDM simulation with the overall results from current photometric and spectroscopic surveys {e.g., SDSS, SEGUE, RAVE) it is far from straightforward to constrain the models with the available data,  mainly for two reasons: 1) the resolution in the models is too low to allow the investigation of the small scale features that we actually see in a real galaxies, 2) the selection of targets in the observational studies is either not well understood/documented or not optimized for the particular question of interest. Future surveys, especially the massive spectroscopic surveys operating after {\\sl Gaia} will be able to have very good selection criteria combined with, almost, exhaustive coverage of a given stellar population, hence enabling a more detailed comparison with models.  The several examples of past studies of the Milky Way and its stellar components given in this review show the need for large spectroscopic surveys to enable a full disentangling of the formation processes involved in shaping our Galaxy. Indeed, with the combination of {\\sl Gaia}'s distances and proper motion and a ground based massive follow-up providing the $6D$ phase-space as well as the multi-dimensional abundance and age space we will finally be able to start testing the models for real \\citep[compare, e.g., the proposed tests in][which are only feasible with this new data]{2009MNRAS.400L..61S}. In addition, prepare to find the unexpected. This means, be ambitious also when it comes to the elemental abundances and aim for abundance errors smaller than 0.05\\,dex (internally) in order to find and tag also the small building blocks  of our Galaxy, i.e. the open clusters.  \\small  %"
    },
    "1107/1107.4014_arXiv.txt": {
        "abstract": "Gravitational waves are predicted by Einstein's theory  of general relativity as well as other theories of gravity. The rotational stability of the fastest pulsars means that timing of an array of these objects can be used to detect and investigate gravitational waves. Simultaneously, however, pulsar timing is used to estimate spin period, period derivative, astrometric, and binary parameters. Here we calculate the effects that a stochastic background of gravitational waves has on pulsar timing parameters through the use of simulations { and}{ data from the millisecond pulsars PSR J0437--4715 and PSR J1713+0747.} We show that the reported timing uncertainties become underestimated with increasing background amplitude { by up to a factor of $\\sim10$ for a stochastic gravitational-wave background amplitude of $A=5\\times 10^{-15}$, { where $A$ is the amplitude of the characteristic strain spectrum at one-year gravitational wave periods}}. { We find evidence for  prominent low-frequency spectral leakage in simulated data sets including a stochastic gravitational-wave background.} We use these simulations along with independent Very Long Baseline Interferometry (VLBI) measurements of parallax to set a 2--sigma upper limit of $A\\le9.1\\times 10^{-14}$.  We find that different supermassive black hole assembly scenarios do not have a significant effect on the calculated upper limits.  We also test the effects that ultralow--frequency (10$^{-12}$--10$^{-9}$ Hz) gravitational waves have on binary pulsar parameter measurements and find that the corruption of these parameters is less than { those due to $10^{-9}$--$10^{-7}$ Hz gravitational waves.} ",
        "introduction": "\\label{intro} Pulsars are highly magnetized, rapidly rotating neutron stars that have spin periods ranging from  seconds to milliseconds. See \\citet{handbook} for a full review of pulsars and their applications. Pulsars can be separated into two broad categories: normal pulsars and millisecond pulsars (MSPs). Normal pulsars are characterized by large periods and period derivatives ($P\\sim 1$ s and $\\dot{P}\\sim 10^{-12}$, respectively) and large magnetic fields ($B\\sim 10^{12}$ G), whereas MSPs have smaller periods and period derivatives ($P\\sim 0.01$ s and $\\dot{P}\\sim 10^{-20}$, respectively) and smaller magnetic fields ($B\\sim 10^{8}$ G). Pulsars are extremely accurate clocks and, specifically, MSPs exhibit much better timing stability as they do not show evidence of rotational instabilities such as timing noise or glitches, which are prevalent in normal pulsars. In fact, many MSPs rival atomic clocks in their fractional stability (see Figure 24 of \\citealt{l08}).  Because of this and the fact that many MSPs are in binary systems, they are extensively used for tests of general relativity. The first evidence for the existence of  gravitational waves (GWs) come from the precise timing of PSR B1913+16 in which the decrease of  the orbital period is in accordance with energy loss due to gravitational radiation, as predicted by general relativity \\citep{tw82}.  Nearly three decades ago, it was first shown that pulsars could be used to directly detect GWs \\citep{saz78,det79} by using the pulsar timing residuals (further discussed in Section \\ref{sim}) to look for a specific GW signature. \\citet{hd83} showed that by exploiting the known correlations that a stochastic GW background (GWB) would induce in the pulsar timing residuals, one can place an upper limit on the GWB. This work introduces what is today known as the Hellings and Downs curve, which is a main feature in many detection algorithms (see, e.g., \\citealt{jhl+05,abc+08}). The concept of a pulsar timing array (PTA) composed of the best timed MSPs was first developed over two decades ago \\citep{r89,fb90}. Today there are three main PTAs in existence with the goal of GW detection using pulsars: the European Pulsar Timing Array (EPTA; \\citealt{jsk+08}), the North American Nanohertz Observatory for Gravitational waves (NANOGrav; \\citealt{jfl+09}), and the Parkes Pulsar Timing Array (PPTA; \\citealt{m08}), all of which are in collaboration to form the International Pulsar Timing Array (IPTA; \\citealt{haa+10}).  \\subsection{Pulsar timing precision and parameter uncertainties} While the high timing precision of MSPs has allowed the precise measurement of many parameters and effects, pulsar timing suffers from the inherent fact that the noise sources that contribute to the timing residuals are not well understood. To first order, our timing residuals are composed of noise that follows Gaussian statistics and that has a white power spectrum. For each time-of-arrival (TOA), the amount of receiver noise is estimated by the TOA uncertainty, as determined from a least-squares fit to a high signal-to-noise ratio template profile. It is, however, common for these uncertainties to be underestimated or for additional sources of white noise to contribute to the timing residuals.  One reason might be self-standarding: in the process of deriving an average TOA for an observation, the integrated pulse profile is correlated with a template profile. If the template profile is generated from addition of many observations, then at some low level the noise from the observation can correlate with the noise in the standard, leading to an underestimation of the TOA uncertainty. Also, the cross-correlation method typically used has been shown to underestimate TOA uncertainties in the high-noise regime \\citep{hbo05}. Poor selection of standard profiles can underestimate the TOA uncertainty as well. All of these effects cause the pulsar timing parameters to be consequently more uncertain than initially expected. Assuming these effects scale with the uncertainties they affect, one can attempt to mitigate this underestimation by multiplication of the TOA uncertainties by the square root of the reduced $\\chi^2$ value of the timing residuals, as is commonly done in other fields.  Underestimation of the white noise contribution is not the only challenge in assessing pulsar timing parameters and their true uncertainties, though. The presence of non-white noise in pulsar timing residuals is very common for young pulsars and, as timing baselines increase, starts to become important in timing of some MSPs as well \\citep{sc10,vbs+08,sns+05}. While the spectral properties and origins of non-white noise in MSPs are as yet not understood because the white noise is still too prominent to allow detailed analysis, it has already been shown that the least-squares fitting performed in pulsar timing reports underestimated uncertainties for the parameters in the timing fit and may even corrupt the parameter values themselves \\citep{vbs+08}.  A technique to improve the least-squares fitting process to mitigate such effects has recently been proposed \\citep{chc+10}.  A third problem may be present in the form of signals that are part of a non-white spectrum -- such as GWs. Not only does this type of signal affect the fitting process as described above, through introduction of power at a very specific frequency it has the potential to be fully absorbed in the timing signature of periodic parameters in the pulsar timing model, such as pulsar position, proper motion, parallax or many binary parameters. On the one hand this can prove useful since independent measurement of parameters such as parallax (by, e.g., VLBI) can be combined with the pulsar timing measurement to place bounds on the amount of corrupting noise present (see, e.g. \\citealt{dvt+08}), but on the other hand it is worthwhile to assess the amount of influence such noise (especially predicted noise sources such as the GWB) can have on pulsar timing parameters.  \\subsection{Pulsar timing and GWs} \\label{sec:ptgw} Low frequency ($10^{-9}$--$10^{-7}$ Hz) GWs are expected from supermassive black hole binary systems (SMBHBs), cosmic strings, and GWs from the big bang and inflationary era of the early universe. GWs from these sources can manifest themselves in different ways. Single nearby SMBHBs can produce  resolvable waves  with periods on the order of years \\citep{wl03}. SMBHBs and cosmic strings can also produce GW bursts \\citep{dv01,smc07,lss09} in which the duration of the GW signal is much less than the observation time. We also expect pulsar timing arrays to be sensitive to a stochastic background of unresolvable sources. This stochastic background can be described by a characteristic strain spectrum $h_c(f)$ defined as a frequency--dependent power law \\citep{jhs+06} \\begin{equation} \\label{eq:strain} h_c(f)=A\\left(\\frac{f}{\\hbox{yr}^{-1}}\\right)^{\\alpha}. \\end{equation} Here $A$ is the dimensionless amplitude and $\\alpha$ is the spectral index. The power in the GWB can then subsequently be written as \\begin{equation} \\label{eq:power} P(f)=\\frac{1}{12\\pi^2}\\frac{1}{f^3}h_c(f)^2. \\end{equation} In accordance with many cosmological theories, one can also write $\\Omega_{{\\rm gw}}(f)$, the energy density per unit logarithmic frequency interval, in terms of the characteristic strain spectrum as \\begin{equation} \\label{eq:omega} \\Omega_{{\\rm gw}}(f)=\\frac{2}{3}\\frac{\\pi^2}{H_0^2}f^2h_c(f)^2, \\end{equation} where $H_0$ is the Hubble constant. This paper will  focus primarily on a stochastic background of SMBHB sources with a spectral index of  $\\alpha = -2/3$ \\citep{wl03,jb03,ein+04}. Recent work has shown that this simple power law does not hold true at higher frequencies ($f> 10^{-8}$) due to the discrete number of sources. At these frequencies the characteristic strain is written as \\citep{svc08} \\begin{equation} \\label{eq:broken} h_c(f)=h_0\\left(\\frac{f}{f_0}\\right)^{\\alpha}\\left(1+\\frac{f}{f_0}\\right)^{\\gamma}, \\end{equation} where $h_0$, $f_0$, and $\\gamma$ are model-dependent parameters based on specific SMBHB merger scenarios. The implications that this result will have on our simulations will be discussed  in Section \\ref{upper} where we test whether this broken power law will make a significant difference to methods that use periodic pulsar parameters with relatively high-frequency fitting functions.   Although a direct detection of the GWB has not been made to date, methods using  pulsar timing residuals have been used to disprove proposed SMBHB systems \\citep{jll+04} and to put limits on SMBHB coalescence rates \\citep{wjy+11} and on the dimensionless strain amplitude \\citep{srt+90,jhs+06,hlj+11}. The most stringent published upper limit to date is  $A\\le 1.1\\times 10^{-14}$  \\citep{jhs+06}.  They use a statistic sensitive to a red spectrum.  In this paper, we present a method which relies on the derived timing \\citep{vbc+09} and VLBI \\citep{dvt+08} parameters  for individual pulsars. It involves adding simulated GWs to pulsar timing parameters and refitting for the pulsar timing parameters to determine the effect of the induced GW.  This idea of using VLBI and timing parameters to place limits on the stochastic background was first introduced in \\citet{dvt+08}. In that work, the kinematic distance obtained from timing measurements of $\\dot{P}_{\\rm b}$ in the presence of a GWB is compared to the VLBI measurement of the distance to place an upper limit.    \\subsection{This paper} \\label{sec:paper} In this paper, we determine how strongly  pulsar parameters can be corrupted by a GWB and red noise in general. We also use independent VLBI measurements to construct upper limits on the stochastic GWB. This is an important question because these corruptions shown here through simulations could be confirmed  by more accurate VLBI measurements as well as other optical interferometry missions such as GAIA \\citep{mk10}. Furthermore, many relativistic tests in the strong field regime could be affected by using the corrupted values of the post-Keplerian parameters $\\dot{P}_{\\rm b}$ and $\\dot{\\omega}$.  This paper is organized as follows: in Section \\ref{sim} we describe the basis of our simulations, including a brief review of GW effects on the pulsar timing residuals and fitting procedures. In Section \\ref{obs} we describe the observations used in our simulations. In Section \\ref{sec:analysis} we present the effect of GWs on pulsar timing parameters and in Section \\ref{discussion} we summarise our findings. ",
        "conclusions": "\\label{discussion} We have introduced a  way to investigate the effects of { low-frequency noise in the form of} a stochastic background of GWs on pulsar timing parameters. { While this work focuses particularly on the stochastic GWB, the methods for investigating the noise/GW induced corruption of pulsar timing parameters could apply to any low-frequency noise source.} This method involves a Monte-Carlo simulation over different possible GWBs, adding them into the pulsar timing residuals, fitting for pulsar timing parameters, and producing distributions of various GW-affected pulsar parameters. These parameter distributions are then analysed to determine the overall effect that the GWB has on pulsar timing. { To summarise our results, we have: \\begin{itemize} \\item Replicated the \\citet{vbs+08} results of red noise leakage and how it affects parameter uncertainties, showing that some pulsar parameters could be underestimated by up to a factor of $\\sim 10$ for a GWB with an amplitude of $A=5\\times 10^{-15}$. \\item Expanded on the above work with another pulsar (PSR J1713+0747), clarifying the importance of ecliptic latitude and the reliability of astrometric parameter measurements. \\item Shown that one can place upper limits on the stochastic GWB by comparing VLBI derived astrometric parameters to those obtained through pulsar timing. \\item Demonstrated practically that the spectral break (or its monochromatic components) are unlikely to affect pulsar timing before the SKA era. \\item Demonstrated that ultra-low frequency GWs have no obvious effect on pulsar timing parameters. \\item Demonstrated that even red noise at very low levels can leak into higher frequency terms and cause parameter corruptions. This lends a partial explanation for EFACs (a multiplicative error factor that is used to normalise the reduced chi--squared in the fitting process), even in a seemingly white data set. \\end{itemize}} Future work could investigate these timing effects on a larger sample of pulsars, including normal (non-MSP) pulsars to check for corruption of pulsar parameters though these effects should be smaller due to larger errors on the timing parameters."
    },
    "1107/1107.6011_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  The origin of the large-scale structures of our universe can be explained by the presence of primordial density perturbations in the early universe, which were of super-horizon scales and almost scale-invariant. However, understanding how such density perturbations were seeded has remained a big mystery for cosmology. On the other hand, cosmic inflation~\\cite{Starobinsky:1980te,Sato:1980yn,Guth:1980zm} has been known to be able to generate super-horizon and (nearly) scale-invariant fluctuations for scalar fields. (Early universe models that can generate such field fluctuations without having an inflationary era have also been studied, see e.g.~\\cite{Mukohyama:2009gg,Rubakov:2009np,Creminelli:2010ba}.) Then now, one may rephrase the question about the origin of the large-scale structures as how such field fluctuations were converted into the primordial density perturbations of the universe, which eventually lead to structure formation.  An attractive mechanism for creating the primordial density perturbations from the field fluctuations is the curvaton mechanism~\\cite{Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct}, in which the curvaton field possessing large-scale field fluctuations generates the density perturbations while it oscillates about its potential minimum and comes close to dominating the universe. The idea of the curvaton mechanism is thus simple, and also has the merit that when embedded in inflationary cosmology, it frees the inflaton from being responsible for generating the perturbations and therefore drastically relaxes constraints on inflationary model building. However, while there have been extensive works on the curvaton mechanism, most of the literature has focused on curvatons possessing rather trivial potentials (e.g. quadratic potentials), thus lead to common predictions for the density perturbations, such as large/small non-Gaussianity in the perturbations depending on whether the curvaton is subdominant/dominant upon decay.  In this paper, we explore density perturbations sourced by a curvaton rolling along an arbitrary energy potential. The key feature of a curvaton potential which deviates from a quadratic one is that the curvaton field with large-scale field fluctuations starts its oscillation at different times at different patches of the universe. Such behaviour contributes to the curvaton energy density perturbations in addition to that sourced directly from the original field fluctuations. We find that such non-trivial conversion processes of the field fluctuations into the density perturbations can lead to strong enhancement/suppression of the density perturbations of the universe, as well as for their non-Gaussianities. In particular, we will show that the non-Gaussianity parameter~$f_{\\mathrm{NL}}$ can become large with either sign no matter if the curvaton dominates or subdominates the universe when it decays. Several previous works have considered specific cases, e.g., curvaton potentials with a non-quadratic polynomial term were studied in~\\cite{Enqvist:2005pg,Enqvist:2008gk,Enqvist:2009zf}, and cosine type potentials in~\\cite{Kawasaki:2008mc} which clearly pointed out that non-uniform onset of curvaton oscillations can have important effects on the density perturbations. Related discussions can be found in~\\cite{Sasaki:2006kq}, though such effects were not explicitly taken into account. In this paper we carry out a systematic study of density perturbations from curvatons with generic potentials through analytic analyses.  It is worthwhile to investigate curvaton potentials which deviate from simple quadratic ones, both from observational and theoretical reasons: Considering a curvaton whose field fluctuations were generated during the inflationary era (which is the case studied in this paper), in order for the curvaton to produce the density perturbation spectrum which is red-tilted (i.e. with negative spectral index~$n_s\\! - \\! 1$) as suggested by latest CMB observations~\\cite{Larson:2010gs,Dunkley:2010ge,Hlozek:2011pc} without relying on specific inflation mechanisms,\\footnote{For a non-tachyonic curvaton (i.e. a curvaton located along a potential with non-negative curvature during inflation), the spectral index of the resulting density perturbation spectrum is set by the time variation of the Hubble parameter during inflation as $ n_s -1 \\geq 2 \\dot{H}/H^2$ (cf.~(\\ref{ns-1})). Hence, assuming single-field canonical slow-roll inflation, the Lyth bound~\\cite{Lyth:1996im} relates the spectral index with the inflaton field~$\\phi$ range as \\begin{equation*} \\frac{1}{M_p} \\left| \\frac{d\\phi}{d\\mathcal{N}} \\right| \\geq \\sqrt{1-n_s}, \\end{equation*} where $M_p$ is the reduced Planck mass and $\\mathcal{N}$ the e-folding number. This shows that unless the curvaton is tachyonic during inflation, the WMAP~\\cite{Larson:2010gs} central value $n_s  \\approx 0.96$ requires a super-Planckian field range for the inflaton.\\label{foot1}} the curvaton needs to be located along a potential with negative curvature during inflation (cf.~(\\ref{ns-1})). This obviously suggests the curvaton potential to take non-trivial forms. Furthermore, explicit curvaton models constructed in the framework of microscopic physics can naturally possess intricate energy potentials, e.g., curvatons in supersymmetric models~\\cite{Hamaguchi:2003dc}, a stringy curvaton model in~\\cite{Kobayashi:2009cm} from a D-brane moving along the internal compactified space giving rise to a multi-dimensional periodic curvaton potential. \\cite{Burgess:2010bz}~also discusses a string realization of the curvaton, where the potential is a linear combination of exponential terms.  We also apply our general results to the case where the curvaton is a pseudo-Nambu-Goldstone (NG) boson~\\cite{Dimopoulos:2003az} of a broken U(1) symmetry. The periodic curvaton potential necessarily possesses not only minima but also maxima, around which the potential curvature is negative so that the resulting density perturbation spectrum can become red-tilted. We explore the parameter space of NG curvatons producing a red-tilted density perturbation spectrum as suggested by observations, and show that the allowed regions for the inflationary energy scale and the reheating temperature (i.e. when the inflaton decays) are strongly constrained, unless the NG curvaton is located close to the maximum of its potential during inflation. In such hilltop cases, the resulting density perturbations are extremely enhanced, thus relaxing the bounds on the inflation/reheating scales. Non-Gaussianity of the density perturbations also increases in the hilltop limit to values $10 \\lesssim f_{\\mathrm{NL}} \\lesssim 30$ in most part of the allowed parameter window, which will be tested by upcoming CMB measurements. Such mild increase of non-Gaussianity towards the hilltop limit is actually a rather generic feature of hilltop curvatons, as will be shown in the paper.  This paper is organized as follows. We derive analytic expressions for density perturbations from curvatons with generic potentials in Section~\\ref{sec:pert}. Then in Section~\\ref{sec:flat}, we move on to discuss special cases where the curvaton starts its oscillation from a flat region of its potential, including hilltop curvatons. It will be shown that in such cases, effects due to fluctuations of the onset of the oscillation become crucial. As a simple example of a curvaton model generating density perturbations that match with observational constraints, we investigate the case where the curvaton is a pseudo-NG boson in Section~\\ref{sec:PNGC}. We present our conclusions in Section~\\ref{sec:conc}. Discussions on dynamics of scalar fields in an expanding universe are provided in Appendix~\\ref{app:SFD}. While we focus on sinusoidal curvaton oscillations in the main body of the paper, in Appendix~\\ref{app:non-sinu} we further generalize our discussions to incorporate non-sinusoidal curvaton oscillations. Curvatons along linear potentials are also discussed in Appendix~\\ref{app:non-sinu}. ",
        "conclusions": "\\label{sec:conc}  In this paper, we have explored density perturbations from a curvaton with a generic potential, and shown that the curvaton's field fluctuations can lead to non-uniform onset of the curvaton oscillation~$\\delta H_{\\mathrm{osc}}$. This sources additional contributions to the resulting density perturbations, in addition to that from the field fluctuations directly perturbing the curvaton energy density. We showed that the combination of the two effects gives rise to non-trivial behaviours for density perturbations from curvatons, such as the curvaton sourcing large $f_{\\mathrm{NL}}$ with either sign even if it dominates the universe before decay. The effect due to $\\delta H_{\\mathrm{osc}}$ becomes important when the curvaton potential deviates from the familiar quadratic type. Since a curvaton potential with negative curvature is required during inflation in order to explain the observationally suggested red-tilted density perturbation spectrum (without invoking inflation with large $|\\dot{H}/H^2|$), it is of importance to evaluate effects from~$\\delta H_{\\mathrm{osc}}$. Our work was also motivated by curvaton models based on microscopic physics such as supersymmetry or string theory, which provides explicit examples of curvaton potentials possessing non-trivial forms.  As a simple case where the $\\delta H_{\\mathrm{osc}}$ effects become dominant on the generated density perturbations, we have studied curvatons that start their oscillations from a flat region along its potential. Especially for hilltop curvatons, we showed the strong enhancement of the linear order density perturbations~$\\mathcal{P}_{\\zeta}$, accompanied by a mild increase of the non-Gaussianity~$f_{\\mathrm{NL}}$.  Another non-trivial case, which we have not studied intensively in this paper, is when various contributions to the density perturbations cancel each other and thus suppress the perturbation amplitude. Conversely, then the non-Gaussianity is expected to become large. It would be interesting to investigate such cases, which may be realized by, e.g., curvaton potentials possessing plateau regions, or with superimposed repeated modulations.\\footnote{Superimposed repeated modulations to the inflaton potential can also have interesting consequences to the density perturbation spectrum generated by the inflaton itself, see e.g.~\\cite{Adams:1992bn,Wang:2002hf,Feng:2003mk,Chen:2008wn,Pahud:2008ae,Hamann:2008yx,Flauger:2009ab,Hannestad:2009yx,Flauger:2010ja,Kobayashi:2010pz}.} We leave this for future work.  We also applied our generic results to study a simple model where the curvaton is a pseudo-Nambu-Goldstone (NG) boson of a broken U(1) symmetry. Although NG curvatons away from the hilltop require high inflation/reheating scales in order to generate red-tilted density perturbation spectra with $ n_s - 1 \\sim - 10^{-2}$, it was shown that NG curvatons at the hilltop can work with a wide range of inflation/reheating scales. We also showed that working NG curvatons in the hilltop predict the non-Gaussianity to lie in the range $10 \\lesssim f_{\\mathrm{NL}} \\lesssim 30$, which is accessible to upcoming CMB observations such as the PLANCK satellite.  For the four-point correlation functions of the density perturbations from hilltop NG curvatons, one obtains relations that are similar to those for the familiar quadratic curvatons, i.e., $\\tau_{\\mathrm{NL}} \\sim g_{\\mathrm{NL}} \\sim f_{\\mathrm{NL}}^2$. However, different curvaton potentials may lead to different behaviours among the non-Gaussianities (see related discussions in~\\cite{Enqvist:2008gk}). It will be interesting to examine systematically the range of possibilities arising from curvatons with generic energy potentials.  Both from observational and theoretical reasons, it is important to investigate in detail the conversion process of the curvaton field fluctuations into the resulting density perturbations. In this work, we have provided generic analytic expressions incorporating effects due to non-uniform onset of the curvaton oscillation, which open up new possibilities for the curvaton paradigm."
    },
    "1107/1107.4825_arXiv.txt": {
        "abstract": "We present calculations of the time evolution of the prompt spectra of gamma-ray burst models involving generic internal dissipation regions, including internal shocks, either by itself or in the presence of an external photon source such as a photosphere. The method uses a newly developed time-dependent code involving synchrotron emission and absorption, inverse Compton scattering and pair formation. The models reproduce the typically observed Band spectra and their generic time evolution, including the appearance of an extra keV-GeV component, whose delay in simple SSC models, however, is only partially able to explain the several seconds observed GeV delays. On the other hand, models involving both a photosphere and an internal dissipation region at a larger radius produce both an extra GeV component and time delays which are in the range of the observations. ",
        "introduction": "\\label{sec:intro}  Most of the prompt emission spectra of gamma-ray bursts (GRBs) can be described by the well-known Band function \\citep{ban93}, which has a spectral peak around the MeV range. Since the EGRET era \\citep{gon03}, distinct spectral components in higher energy ranges have been reported. Moreover, bright optical emissions during the prompt phase, such as in GRB 080319B \\citep{rac08}, have also been detected. Such lower energy emissions also consist of distinct spectral components. In the popular internal shock model of GRB the prompt MeV gamma-rays are attributed to synchrotron radiation from electrons accelerated in shocks within the GRB outflow \\citep{pir05,mes06}, while higher energy photons are attributed to inverse Compton (IC) upscattering of these synchrotron photons, a scheme called synchrotron self-Compton or SSC \\citep{mes94,pil98,pan00,gue03}. Variants on this scheme involve having the MeV photons, whether of synchrotron or other origin, arising at a different location (usually at a smaller radius) than the location where they are upscattered, a scheme refereed to as external inverse Compton or EIC \\citep{bel05,tom09,mur11}. The generic radiation physics is independent, to a large degree, of the specific macroscopic model for the acceleration of the radiating or scattering non-thermal particles. Thus, the results we discuss here are not tied specifically to internal shock models, but apply also to more generic dissipative models, in which synchrotron, IC and, in the case of high radiation density also pair formation, are the main radiation mechanisms.  The recent observations over a wide range of photon energies by the {\\it Fermi}-Large Area Telescope (LAT) and GLAST Burst Monitor (GBM) instruments and robotic optical telescopes such as TORTORA have stimulated research on the temporal evolution of the GRB spectra, which is a key for distinguishing various models. Several authors have developed time-dependent codes and discuss the spectral evolution in GRBs \\citep{pee05,pee08,bel08,vur09,bos09,dai11}. Some of those studies focus on the effects of the dynamical evolution of the shocked region, multiple Compton scattering, or thermal particles etc. An expanding range of physical situations needs to be considered, in view of the ongoing observational and theoretical progress being made. Here, we improve our previous Monte Carlo code, developed in a series of studies for hadronic and leptonic processes \\citep[see, e.g.,][]{asa07}, into a time-dependent code. Our code is based on the one-zone approximation with the Lagrangian specification in the energy space for particles. Although we plan to develop this code to deal with the hadronic processes in the future, at this time only the leptonic processes are included, as a first step.  Our goal in this paper is to study the temporal behavior of the radiation field in the typical types of geometries generally considered in GRB models. In particular, we study the development of the photon spectra as a function of time, as well as the expected light curves in specific energy ranges, in order to gain insight into  current questions such as the origin of the delay between MeV and GeV pulses in bursts observed by the {\\it Fermi} \\citep{916C,902B,926A}, or the presence of additional spectral components above or below \\citep{510,902B,926A} the usual Band (MeV) spectra known from previous Compton and {\\it Swift} studies.  In \\S \\ref{sec:model} we present the basic physical model used in our numerical, discussing the various physical processes involved and the way the time-dependent spectra and light curves are computed. The results are presented in \\S \\ref{sec:results}. In subsection \\S \\ref{sec:moderate} we discuss first three cases characterized by parameters such as inferred from pre-Fermi GRB observations, in the context of an SSC model. These ``moderate'' models serve as a reference point, providing insight into what was missed in the pre-Fermi observations.  In subsection \\S \\ref{sec:lat} we explore the temporal and spectral properties and predictions for bursts with parameters which might characterize some of the more interesting bursts observed by the Fermi LAT instrument, again based on an SSC model. In \\S \\ref{sec:external} we discuss the properties of a model including external photon sources providing a Band function spectrum which is reprocessed and upscattered at a different location.  We summarize our numerical results in \\S \\ref{sec:sum} and discuss their relevance and implications in \\S \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:sum}  The results of our time-dependent simulations for the internal dissipation model are summarized as follows. \\begin{itemize} \\item The temporal evolution of the photon energy distribution inside the shell affects the resultant spectra seen by the observer. Especially for weak magnetic fields, the evolution of the IC component shows the limitation on the steady state approximation. Time-resolved spectra are necessary in order to verify expected features of the spectral evolution, such as the decrease of the cut-off energy due to $\\gamma \\gamma$-absorption, the growth of the IC component, and a spectral softening due to electron cooling. \\item Even in our one-zone approximation, the power-law spectrum is produced above the $\\gamma \\gamma$ cut-off energy. The spectral evolution and photon escape during the electron injection are essential for this spectral shape. \\item We confirm that the light curves follow the FRED shape \\citep{fen96}. The pulse widths and peak times are affected by not only the curvature effect but also the intrinsic evolution of the energy distribution. The decrease of the cut-off energy leads to an earlier GeV peak time, and the gradual electron cooling causes the delayed onset and broad pulse profile of the optical emission. \\item The possible deviation mechanisms from the simple Band function (other than IC emission) are synchrotron emissions from secondary electron-positron pairs, electron heating via SSA, and late synchrotron emission after the injection ends. The latter two mechanisms yield a spectral bump in the low energy portion, while the secondary pairs enhance the synchrotron flux in the high energy range. \\item The retarded growth of the IC component can cause the delayed onset of GeV emission. This may be a partial reason for the delayed onsets seen in the {\\it Fermi}-LAT GRBs. \\item The combination of the IC and late synchrotron components can produce the ``extra spectral component'' from keV to GeV as seen in the {\\it Fermi}-LAT GRBs. This is a counter example to the claim that the IC emission cannot explain the extra components in the {\\it Fermi}-LAT GRBs. \\item When the IC and late synchrotron components are prominent, a spectral evolution from the Band function to a power-law dominated shape is seen. This is similar to the spectral evolution observed in GRB 090510. \\item We confirm the essential features of external photon models such as proposed by \\citet{tom09,tom10}. The IC emission using external Band photons coming from inside regions can reproduce the delayed onset of GeV emission. The effect of the anisotropic emissivity adds an extra factor for the delay timescale. The evolution of the spectral indices is similar to those seen in GRB 080916C. The long-lasting tail of GeV emission in this model is also an interesting feature for the {\\it Fermi}-LAT GRBs. \\end{itemize}"
    },
    "1107/1107.1237.txt": {
        "abstract": "We have assembled the largest sample of ultra hard X-ray selected (14-195~keV) AGN with host galaxy optical data to date, with 185 nearby (z$<$0.05), moderate luminosity AGN from the ${\\sl Swift}$ BAT sample.  The BAT AGN host galaxies have intermediate optical colors ($u-r$ and $g-r$) that are bluer than a comparison sample of inactive galaxies and optically selected AGN from the Sloan Digital Sky Survey (SDSS) which are chosen to have the same stellar mass.  Based on morphological classifications from the RC3 and the Galaxy Zoo, the bluer colors of BAT AGN are mainly due to a higher fraction of mergers and massive spirals than in the comparison samples.  BAT AGN in massive galaxies (log M$_*$$>$10.5) have a 5 to 10 times higher rate of spiral morphologies than in SDSS AGN or inactive galaxies.  We also see enhanced far-IR emission in BAT AGN suggestive of higher levels of star formation compared to the comparison samples. BAT AGN are preferentially found in the most massive host galaxies with high concentration indexes indicative of large bulge-to-disk ratios and large supermassive black holes.  The narrow-line (NL) BAT AGN have similar intrinsic luminosities as the SDSS NL Seyferts based on measurements of  [O \\textsc{III}] $\\lambda5007$.   There is also a correlation between the stellar mass and X-ray emission. The BAT AGN in mergers have bluer colors and greater ultra hard X-ray emission compared to the BAT sample as whole. In agreement with the Unified Model of AGN, and the relatively unbiased nature of the BAT sources, the host galaxy colors and morphologies are independent of measures of obscuration such as X-ray column density or Seyfert type.  The high fraction of massive spiral galaxies and galaxy mergers in BAT AGN suggest that host galaxy morphology is related to the activation and fueling of local AGN. ",
        "introduction": "Most galaxies with bulges harbor a supermassive black hole in their center \\citep{Magorrian:1998p9015}, yet only a small fraction exhibit the powerful radiative or kinetic output associated with an active galactic nucleus (AGN).  While it is well established that matter falling onto the supermassive black hole is emitted as energy, the source of this material remains highly controversial.  To understand what activates and continues to fuel AGN, we must better characterize the conditions of the host galaxies in which they are found.  What environmental factors trigger these black holes to begin emitting so much energy and what continues to fuel this process?  Numerical simulations suggest that quasars (L$_{\\mathrm{bol}}$$>$$10^{45}$ erg s$^{-1}$) are the end product of mergers between gas-rich disk galaxies, and that supermassive black hole accretion heats the interstellar material and quenches star formation leading to passive elliptical galaxies \\citep{DiMatteo:2005p5934}.  Alternatively, other simulations suggest sources other than mergers may fuel lower luminosity AGN, such as gas streaming down galactic bars or steady cold gas streams \\citep{Mulchaey:1997p9021,Hopkins:2006p9574,Dekel:2009p9116}.  A number of observational studies have provided interesting yet contradictory results about the relationship between the host galaxy and the AGN.  A study of the host galaxies of X-ray selected AGN from the Extended Chandra Deep Field-South found that AGN are in the most luminous galaxies, with intermediate optical colors, and bulge dominated morphologies \\citep{Silverman:2008p8462}.  Another study of narrow emission line (NL) AGN in the Sloan Digital Sky Survey (SDSS) found the hosts were predominantly massive early-type galaxies and the most luminous AGN galaxies had significant star formation \\citep{Kauffmann:2003p2397}.  An additional survey of the SDSS NL AGN host galaxies found, compared to a sample of nearby inactive galaxies, most AGN occur along the red sequence  \\citep{Westoby:2007p8467}.  Even though these studies draw their conclusions from large optical surveys or soft X-ray surveys, their results may be biased by missing an important population of obscured AGN.  A less biased survey of AGN must account for absorption of light by gas and dust.  In addition, the survey must distinguish between host galaxy emission primarily from stars and emission from AGN.  Historically, an inability to account for these factors provided different results depending upon the wavelength used for observation. Initially, AGN were selected by radio techniques \\citep{Heckman:1980p9959} or optically using broad and strong emission lines, a luminous point-like nucleus, or irregular galaxy color \\citep{Weedman:1977p8832}.  The presence of broad optical emission lines was used to separate Seyfert 1 galaxies with broad permitted and narrow forbidden lines and Seyfert 2 galaxies with only narrow permitted and forbidden line emission.  In the unified model (Antonucci 1993) both broad line Seyfert 1 galaxies and obscured Seyfert 2 galaxies are intrinsically the same, with the differences being the viewing of the central engine.  Surveys of AGN taken in the optical, UV, and soft X-rays ($<$ 2 keV) miss an important population of obscured narrow line AGN only visible in the ultra hard X-ray and mid-IR wavelengths \\citep{Mushotzky:2004p8576,Koss:2011p12483}.  While the mid-IR wavelength is less obscured, this wavelength range is problematic because of confusion with emission from star formation, sensitivity to the amount of obscuring material, and the lack of a unique way to select AGN from other luminous IR galaxies \\citep{Stern:2005p8982,Hickox:2009p707}.  Therefore, the ultra hard X-ray, $>$15 keV range offers an important new way to select AGN for a less biased survey.  The BAT survey is an all sky survey in the ultra hard X-ray range that has identified 461 objects of which 262 are AGN \\citep{Tueller:2010p6018}.  The BAT instrument is a large field of view (1.4 steradian) coded aperture imaging instrument.  Because of the large position error of BAT ($\\approx2\\arcmin$) higher angular resolution X-ray data for every source from Swift-XRT or archival data have been obtained allowing associations with 97\\% of BAT sources.  This sample is particularly powerful since the BAT is sensitive in the 14--195 keV band and at obscuring columns of $>$$10^{24}$ cm$^{-2}$ where only high-energy X-ray emission (tens of keV) can pass through the obscuring material.  It is therefore sensitive to heavily obscured objects where even hard X-ray surveys (L$_{2-10 \\: \\mathrm{keV}}$) are severely reduced in sensitivity.    At 22 months\\footnote{http://heasarc.gsfc.nasa.gov/docs/swift/results/bs22mon/}, the BAT survey has a sensitivity of approximately $2.2\\times10^{-11}$ erg cm$^{-2}$ s$^{-1}$.  With this sensitivity, the BAT survey is about 10 times more sensitive than the previous all-sky ultra hard X-ray survey, HEAO 1 \\citep{Levine:1984p9096}.   About 15$\\%$, or 30 of the AGN, have never before been detected as AGN at other wavelengths.  While AGN host galaxy studies using X-ray surveys typically probe objects at moderate redshift (out to z$\\approx$1), more nearby (z$<$0.05) AGN offer the best opportunity to study the host in detail since high spatial resolution data are easily obtainable. The majority of the BAT AGN are nearby with a median redshift of 0.03.  Thus, a study of the BAT AGN sample provides an excellent opportunity to answer the controversial question of AGN fueling and its relationship to the host galaxy.  %The multiwavelength properties of the BAT AGN sample have been studied in a number of papers ().  The BAT AGN sample host galaxies have been studied in optical imaging by Schawinski et al.~2009 and Vasudevan et al.~2009.  However, these surveys were limited to 16 and 17 BAT AGN whereas we focus on a much larger sample of 185 BAT AGN.  In addition, these papers reached opposite conclusions, the Schawinski paper found AGN seeming to be in redder galaxies indicating the AGN may be suppressing star formation and the Vasudevan study found AGN in bluer galaxies.  With 10 times more AGN than the previous optical imaging studies, we have a much larger sample to study the host galaxies. ",
        "conclusions": "We have assembled the largest sample of ultra hard X-ray selected AGN with host galaxy optical data to date, with 185 AGN in total.  We have performed extensive modeling with GALFIT to effectively remove the AGN light from the optical images.  Using optical photometry, morphology, and spectroscopy, along with FIR emission we found:   \\begin{enumerate} \\renewcommand{\\theenumi}{(\\roman{enumi})} \\renewcommand{\\labelenumi}{\\theenumi} \\item The BAT AGN galaxies are bluer in optical color than inactive galaxies or SDSS Seyferts of the same stellar mass. \\item We find a much higher incidence of spiral morphologies in BAT AGN compared to SDSS AGN or inactive galaxies.  Amongst massive galaxies ($\\log$ M$_*>10.5$), the BAT AGN show a preference for spiral morphologies that is 5 to 10 times higher than SDSS AGN or inactive galaxies.  We also find that the bluer colors of BAT AGN can be accounted for by a higher fraction of mergers and spirals. \\item The BAT AGN have greatly enhanced 90 $\\mu$m emission compared to inactive galaxies or SDSS Seyferts matched in redshift and stellar mass. \\item The BAT NL AGN have similar intrinsic [O \\textsc{III}] $\\lambda5007$ luminosities as NL SDSS Seyferts of the same redshift range. \\item The BAT AGN are found in the most massive host galaxies with high concentration indexes indicative of large bulge-to-disk ratios and large supermassive black holes. \\item We also find that the average ultra hard X-ray luminosity increases with stellar mass and that BAT AGN in mergers have greater ultra hard X-ray emission than those in other morphological types.  This suggests a link between supermassive black hole growth and the mass of the host galaxy. \\item In agreement with the Unified Model of AGN, we find the host galaxy colors and morphology are independent of X-ray column density and optical Seyfert classification. \\end{enumerate} %on Galaxy Zoo and the RC3.  We also confirm the increased number of spirals based on the higher galaxy inclinations of BAT AGN.  Finally, we confirm the smaller number of ellipticals in the BAT AGN sample using SDSS photometry and spectroscopy catalog parameters  These results indicate that host galaxy morphology is related to the activation and fueling of local AGN.  Ultra hard X-ray selected AGN are particularly associated with massive spiral galaxies and galaxy mergers.  These types of objects are generally associated with bluer colors, compared to the red massive early-type galaxies at similar stellar masses.  These observational results provide some evidence for an association between AGN activity and galaxy mergers (e.g., di Matteo et al.~2005), and also provide examples of AGN activity driven by the stochastic accretion of cold gas that should be more prominent among late-type systems \\citep{Hopkins:2006p9574}.  Recent simulations have also suggested a transition between the fueling mechanisms of AGN with nonmerger events predominantly powering lower luminosity AGN and merger-induced fueling dominant in more luminous quasars  \\citep{Hopkins:2009p10071}.  We may be seeing evidence of this transition in our sample of BAT AGN that is powered both through merger events and less powerful nonmergers such as accretion of cold gas in late type systems.  In support of this, we find that BAT AGN in mergers have a greater ultra hard X-ray emission than those in other morphological types.  However, only a very small fraction (5/185) of BAT AGN in this sample are above the minimum bolometric luminosity associated with quasars (L$_{\\mathrm{bol}}$$>$$10^{45}$ erg s$^{-1}$).  These results suggest that the process of merging may be important for powering more moderate luminosity AGN as well (see also Koss et al.~2010).  In interpreting the results of an X-ray flux limited survey, it is useful to remember that the observed flux is a product of the black hole mass and accretion rate.  On average, more massive galaxies will tend to have higher mass black holes that will produce a larger average X-ray flux than smaller galaxies with on average smaller black holes.  However, among massive galaxies, elliptical morphologies are much more common than spirals, yet we find the most luminous hard X-ray AGN almost exclusively in spiral morphologies.  This suggests that spiral morphologies must have higher accretion rates than elliptical morphologies.  This finding is in agreement with recent theoretical predictions that suggest that only spirals typically have enough gas to trigger higher levels of radiatively efficient accretion in a geometrically thin disk \\citep{Fanidakis:2011p10087}.  In order to understand this further, we are in the process of accurately measuring black hole masses to study the accretion rates for this sample.  Previous optical surveys have found that AGN tend to be in massive galaxies \\citep{Kauffmann:2003p2397}, occur along the red sequence \\citep{Westoby:2007p8467}, and tend to have similar numbers of galaxy mergers as inactive galaxies \\citep{Li:2006p5063}.  However, in an ultra hard X-ray survey of AGN, we find that AGN host galaxies are bluer than inactive galaxies with higher numbers of massive spirals and galaxy mergers.   We do not find observational evidence that the AGN suppresses star formation.  It is surprising that the optical morphologies and colors of ultra hard X-ray selected AGN are so different than emission line selected Type 2 Seyferts given their similar bolometric luminosity as measured in [O III].  However, these results are consistent with recent Spitzer surveys that have found that the AGN detection rate in late-type galaxies and mergers is much larger than what optical spectroscopic observations suggest \\citep{Satyapal:2008p9525,Goulding:2009p6170, Veilleux:2009p7250}.  Finally, studies of X-ray selected AGN at higher redshifts, have also found a significant population of AGN classified as star forming using emission line diagnostics (Yan et al.~2011, accepted).  In the BAT AGN sample, there are several results that suggest optical emission line classification may be biased against late-type galaxies and mergers. In this study, we found that the axis ratios of BAT AGN are in general more inclined and have greater levels of internal extinction than comparable SDSS AGN.  This extinction could obscure or dilute the narrow-line region and cause AGN galaxies to be misclassified as star forming regions.  This finding is also in agreement with a previous analysis of BAT AGN that found optical emission line diagnostics preferentially misclassify merging AGN because of optical extinction and dilution by star formation \\citep{Koss:2010p7366}.  Another possibility is that the BAT AGN may be much more intrinsically luminous than their [O III] emission suggests.  Since the majority of BAT AGN either have broad lines or are NL AGN that are correctly classified as Seyferts, yet are found to have much greater hard X-ray luminosities, this must be an important factor.  In support of this, two studies of BAT AGN have found a very weak correlation between the [O III] and hard X-ray luminosity and that BAT AGN have additional reddening of the narrow line region not accounted for in optical studies \\citep{Winter:2010p6825,Melendez:2008p3807}.  This is also supported by the much greater number of narrow-line SDSS Seyferts compared to hard X-ray selected AGN.  In the SDSS survey area, there are 24 optical emission line selected narrow-line Seyferts detected for each hard X-ray AGN at the same redshift.  Some of these undetected sources may be heavily absorbed Compton-Thick AGN missed in the hard X-rays, but even the highest estimates expect only $\\approx50\\%$ of local narrow-line AGN are Compton~Thick \\citep{Risaliti:1999p12126}. If the BAT AGN are intrinsically more luminous than emission line selected AGN, this may explain their higher rates of mergers and enhanced FIR emission.    We are currently in the process of assembling a larger survey of optical spectra of BAT AGN to better understand optical and X-ray measures of intrinsic luminosity."
    },
    "1107/1107.1777_arXiv.txt": {
        "abstract": "We present a comprehensive characterization of the general properties (luminosity functions, mass, size, ages, etc) of optically selected compact stellar objects (knots) in a representative sample of 32 low-z Luminous and Ultraluminouos Infrared Galaxies, (U)LIRGs. It is important to understand the formation and evolution of these properties in such systems, which represent the most extreme cases of starbursts in the low-z Universe. We have made use of high angular resolution ACS images from the Hubble Space Telescope in \\filterb ($\\sim$ \\textit{B}) and \\filteri ($\\sim$ \\textit{I}) bands. The galaxies in the sample represent different interaction phases (first contact, pre-merger, merger and post-merger) and cover a wide luminosity range (11.46 $\\leq$ log (\\lir/\\lsun\\twospace) $\\leq$ 12.54). With a median size of 32 pc, most of the nearly 3000 knots detected consists of complexes of star clusters. Some of the knots ($\\sim$15\\%) are so blue that their colors indicate a young (i.e., $<$ 30 Myr) and almost extinction-free population. There is a clear correlation of the mass of these blue knots with their radius, where $M \\propto R^{1.91\\pm0.14}$, similar to that found in complexes of clusters in M51 and in giant molecular clouds.  This suggests that the star formation within the knots is proportional to the gas density at any given radius. This relation does not depend significantly on either the infrared luminosity of the system or on the interaction phase. The star formation of all the knots is characterized by luminosity functions (LFs) of the knots with slopes close to 2. Nevertheless, we see a marginally significant indication that the LF evolves with the interaction process, becoming steeper from early to advanced merger phases. Due to size-of-sample effects we are probably sampling knots in ULIRGs intrinsically more luminous (by a factor of about four) than in less luminous systems. They also have sizes and are likely to have masses characteristic of clumps in galaxies at z$\\gtrsim$1. Knots in post-mergers are on average larger ($\\times$ 1.3-2), more luminous (2 mag) in the \\textit{I} band, and 0.5 mag redder than those in systems in earlier phases. Two scenarios are briefly discussed: (1) the likely presence of relatively high extinction in the most advanced mergers; (2) the dissolution of the less massive clusters and/or their coalescence into more massive, evolved superclusters. ",
        "introduction": "Star clusters tend to form where strong star formation occurs and, specially, in starbursts triggered by galaxy interactions and mergers (\\citealt{Schweizer98}). During the last two decades young and massive compact clusters  have been extensively studied in different environments including starburst galaxies~\\citep{Meurer95a}, barred galaxies~\\citep{Barth95}, spiral disks~\\citep{Larsen99a}, interacting galaxies (\\citealt{Whitmore99,Whitmore07};~\\citealt{Bik03};~\\citealt{Weilbacher00,Weilbacher03};~\\citealt{Knierman03};~\\citealt{Peterson09};~\\citealt{Mullan11}) and compact groups of galaxies (\\citealt{Iglesias-Paramo01};~\\citealt{Gallagher10};~\\citealt{Konstantopoulos10}).  With masses of $10^{4}-10^{6}$ \\msun and a median size of 3.5 pc (in terms of \\reff), the light of these clusters follow the same luminosity function (LF) in the optical in a wide range of galactic environments, a power-law ($\\psi$(L)dL$\\propto$L$^{-\\alpha}$dL) with an index of $\\alpha$ = 2. Yet, there is some controversy, since the range is rather large depending on the study (\\mbox{1.8 $\\lesssim~\\alpha~\\lesssim$ 2.8}). Recent studies suggest that the LF can be curved, showing a systematic decrease of the slope from roughly 1.8 at low luminosities to roughly 2.8 at high luminosities (\\citealt{Gieles10}). Such truncation can be reproduced by modeling LFs using an underlying cluster IMF with a Schechter function, shape of the mass function (MF) claimed to be observed in star-forming environments (e.g.,~\\citealt{Haas08,Larsen09,Gieles09}). On the other hand, there are works that still favor the universality of the LF and the MF (e.g.,~\\citealt{Whitmore07,Whitmore10};~\\citealt{Fall09};~\\citealt{Chandar11}).   These massive clusters do not form normally in isolation but tend to be clustered themselves (\\citealt{Zhang01,Larsen04}). Star-forming complexes represent the largest units of star formation in a galaxy (see the review by~\\citealt{Efremov95}). They are typically kpc-scale regions encompassing several clusters. Their properties have also been investigated in some nearby interacting galaxies, such as in M51 (\\citealt{Bastian05a}) and in NGC 4038/4039 (\\citealt{Whitmore99}), in order to understand the hierarchy of the star formation in embedded groupings. These works have confirmed that the properties of complexes of clusters resemble those of giant molecular clouds (GMCs) from which they are formed. In particular, the mass-radius relation, namely \\mbox{$M_{GMC} \\propto R_{GMC}^2$}, reflects the state of virial equilibrium in these clouds (\\citealt{Solomon87}) and holds down to the scale of cloud clumps a few parsecs in radius (\\citealt{Williams95}). Complexes of young clusters in M51 have a similar correlation, showing the imprint from the parent GMC (\\citealt{Bastian05a}). Likewise, young massive clusters with masses above 10$^6$ \\msun follow the same relation (\\citealt{Kissler-Patig06}). However, young clusters with lower mass does not (\\citealt{Larsen04,Bastian05b,Kissler-Patig06}). A star formation efficiency that depends on the binding energy (predicted by~\\citealt{Elmegreen97}) of the progenitor GMC could destroy such a relation during the formation of young clusters. Another possible explanation for this lack of a mass-radius relation in young clusters is that dynamical encounters between young clusters (and gas clouds) add energy into the forming clusters, thereby increasing their radii (\\citealt{Bastian05a}).   Although much effort has been made to understand the star formation and evolutionary processes in interacting systems, little is known for Luminous (LIRGs) and  Ultraluminous (ULIRGs) Infrared Galaxies, which represent the most extreme cases of starbursts and interactions in our nearby Universe. LIRGs and ULIRGs are objects with infrared luminosities of $10^{11} L_{\\odot} \\leq L_{bol} \\sim L_{IR}[8 - 1000 \\mu m]$\\footnote{For simplicity, we will identify the infrared luminosity as \\mbox{\\lir(\\lsun) $\\equiv$ log ($L_{IR}[8 - 1000 \\mu m]$) }.} $< 10^{12} L_{\\odot}$ and  $10^{12} L_{\\odot} \\leq L_{IR}[8 - 1000 \\mu m] < 10^{13} L_{\\odot}$, respectively \\citep{Sanders96}. The main source of this luminosity is thought to be the starburst activity, although an AGN may also be present, and even be the dominant energy source in a small percentage of the most luminous systems \\citep{Genzel98a,Farrah03,Yuan10}. They are rich in gas and dust, and more than 50\\% of the LIRGs and most ULIRGs show signs of being involved in a major interaction/merger  \\citep[e.g][]{Surace98,Surace00,Cui01,Farrah01,Bushouse02,Evans02,Veilleux06}.  (U)LIRGs are therefore natural laboratories for probing how star formation is affected by major rearrangements in the structure and kinematics of galactic disks. Establishing the general properties (luminosity functions, mass, size, ages, etc) of the compact stellar objects in (U)LIRGs as a function of luminosity and interaction phase would provide relevant information in order to understand the mechanisms that govern the star formation and evolution in these systems. It is not known either whether the processes that (U)LIRGs undergo produce stellar objects similar or very different in terms of size or luminosity than in less luminous \\mbox{(non-)} interacting galaxies.  In a cosmological context, (U)LIRGs become very relevant as tracers of the infrared phase (i.e., the dust-enshrouded star-forming phase), occurring during the early stages in the formation of massive galaxies. The so-called sub-millimeter galaxies, detected at \\mbox{z $>$ 2}, are similar to the ULIRGs observed in the local Universe, in the sense that they are merging systems with extremely high rates of star formation \\citep{Chapman03,Frayer03,Engel10}. Besides, contrary to what it happens in the local Universe, (U)LIRGs are major contributors to the star formation rate density at z$\\sim$1-2~\\citep{Perez-Gonzalez05}. Star-forming galaxies become increasingly irregular at higher redshifts, with a blue clumpy structure, asymmetry and lack of central concentration (\\citealt{Abraham96,Vandenbergh96,Im99}). It has been claimed that clumps in clumpy galaxies represent star-forming complexes intrinsically more massive by one or two orders of magnitudes than similar-size complexes in local galaxies (\\citealt{Efremov95,Elmegreen09,Gallagher10}). Although most of the clumpy galaxies, observed at \\mbox{z = 0.7-2}, do not show signs of interactions, given the extreme star formation properties in local (U)LIRGs it is relevant to study the properties of their clumpy structure in order to establish the similarities to and differences from high-z clumpy galaxies.   Previous studies of compact stellar structures in (U)LIRGs were focused only on small samples, mostly with low angular resolution ground-based imaging (\\citealt{Surace98};~\\citealt{Surace00}), or on detailed multi-frequency studies of compact stellar objects in individual galaxies (\\citealt{Surace00};~\\citealt{Diaz-Santos07}). Thus, no study has been done so far on a representative sample of luminous infrared galaxies covering the different phases of the interaction process as well as the entire LIRG and ULIRG luminosity range.  This paper presents the first attempt at obtaining an homogeneous and statistically significant study of the photometric properties (magnitudes, colors and sizes) of optically-selected compact stellar objects found in these systems as a function of $L_{IR}$, morphology (i.e., interaction phase) and radial distance to the nucleus of the galaxy. These stellar objects, some of them identified as YMCs candidates for the closest galaxies (i.e., at \\ld$\\lesssim$60 Mpc) and practically all of them identified as cluster complexes further away, will be referred to as ''knots''. The luminosity function of the knots is also evaluated and compared with that obtained in other studies. This study also represents the starting point of a more quantitative analysis based on the direct comparison of these measurements with state-of-the-art simulations of galaxy encounters with linear resolutions comparable to those presented here (e.g., simulations by~\\citealt{Bournaud08a}). The results of this analysis, in particular the evolution of the luminosity functions of the knots, will be presented in a forthcoming paper (Miralles-Caballero et al. 2011, in preparation).  The paper is organized as follows. We first describe the sample and classify the systems according to their morphology in section~\\ref{sec:sample}. The data are presented in section~\\ref{sec:data}. We then explain the analysis methodology, including the source selection, photometric measurements, completeness tests, etc. In section~\\ref{sec:results}  we present the main photometric results, including the properties of the objects measured as a function of \\lir and the interaction phase of the system.  Age and mass derivation for the young objects is also discussed. Finally, the summary and conclusions are given in section~\\ref{sec:summary}. Throughout this paper we use $\\Omega_{\\Lambda}$ = 0.73, $\\Omega_{M}$ = 0.27 and $H_{0}$ = 73 $km s^{-1} Mpc^{-1}$. ",
        "conclusions": "\\label{sec:summary}  A comprehensive study of bright compact knots (most of them apparent associations of clusters) in a representative sample of 32 (U)LIRGs as a function of infrared luminosity and interaction phase has been performed using \\hst\\twospace-ACS \\mbox{B- and} I-band imaging, providing linear resolutions of 10 to 40 pc, depending on the distance of the galaxy. The sample has been divided into three infrared luminosity intervals (low, \\mbox{\\lir $<$ 11.65}; intermediate, \\mbox{11.65 $\\leq$ \\lir $<$ 12.0}; and high, \\mbox{\\lir $\\geq$ 12.0}) and in 4 interaction phases, based on the projected morphology (first approach, pre-merger, merger, post-merger). We have extended toward lower infrared luminosities previous stu\\-dies focused only on ULIRGs (e.g.,~\\citealt{Surace98};~\\citealt{Surace00}), and have detected close to 3000 knots, larger than a factor of ten more than in the previous studies of ULIRGs. The main conclusions derived from this study are the following:  \\begin{enumerate}[-]  \\item The knots span a wide range in magnitudes \\linebreak  (\\mbox{-20 $\\lesssim$ \\mi $\\lesssim$ -9} and \\mbox{-19.5 $\\lesssim$ \\mb $\\lesssim$ -7.5}), and color \\mbox{(-1 $\\lesssim$ \\mbi $\\lesssim$ 5)}. The median values are \\mbox{\\mb=-10.84}, \\mbox{\\mi=-11.96} and \\mbi = 1.0, respectively. \\\\  \\item The knots are in general compact, with a median \\reff of 32 pc, being 12\\% unresolved and a few very extended, up to 200-400 pc. With these sizes, in general (and particularly for galaxies located at 100 Mpc or more) the knots constitute complexes or aggregates of star clusters.  Within the resolution limits there is no evidence of a luminosity dependence on the size. A slight dependence on the interaction phase, in particular in the post-merger phase, needs confirmation with larger samples and better angular resolution.\\\\  \\item A non-negligible fraction (15\\%) of knots are blue and luminous. Given their color (\\mbox{\\mbi$<$0.5}) and luminosity (\\mbox{$<$\\mb\\twospace$>$=-11.5}) they are likely to be young (ages of about 5 to 30 Myr), almost free of extinction, and with masses similar to and up to one order of magnitude higher than the Young Massive Clusters observed in other less luminous interacting systems. \\\\  \\item Unlike for isolated young star clusters in other systems, we find a clear correlation between the knot mass and radius, M$\\propto$R$^{2}$, similar to that found for complexes of star clusters in less luminous interacting galaxies (e.g., M51 and the Antennae) and Galactic and extragalactic giant molecular clouds. This relation does not seem to be dependent on the infrared luminosity of the system or on the interaction phase.  \\item The star formation is characterized by B- and I-band luminosity functions (LFs) with slopes close to 2, extending therefore the universality of the LF measured in interacting galaxies at least for nearby systems (i.e., \\ld$<$100 Mpc), regardless of the total luminosity (i.e., the strength of the global star formation). This result has to be taken with caution because the LF of our knots (which can be mainly formed of complexes of star clusters) may not reflect the same physics as the LF of individual star clusters. Nevertheless, there are slight indications that the LF evolves as the interaction progresses becoming steeper (from about 1.5 to 2 for the I-band) from first approach to merger and post-merger phases.\\\\  \\item Taking into account distance effects, knots in high luminosity systems (ULIRGs) are intrinsically more luminous (a factor of about 4) than knots in less luminous systems. Given the star formation rate in ULIRGs is higher than in less luminous systems, size-of-sample effects are likely to be the natural explanation for this. \\\\  \\item Knots in ULIRGs have both sizes and masses characteristic of stellar complexes or clumps detected in galaxies at high redshifts (z$\\gtrsim$1), as long as their population is about 50 Myr or older. Thus, there is evidence that the larger-scale star formation structures are reminiscent of those seen during the epoch of morphological galaxy transformations in dense environments.  \\item Knots in systems undergoing the post-merger phase are on average redder and more luminous than those in systems in earlier phases. Scenarios involving either a different spatial distribution of the obscuration or a different evolutionary phase in which small knots do not survive for long are considered. The evolutionary scenario could also explain the trend in size toward higher values in post-merger systems. A comparison with numerical simulations to further investigate these scenarios will be presented in a forthcoming paper (Miralles-Caballero et al. 2011, in preparation). \\\\   \\end{enumerate}"
    },
    "1107/1107.2774_arXiv.txt": {
        "abstract": "{The migration of low-mass planets, or type I migration, is driven by the differential Lindblad torque and the corotation torque in non-magnetic viscous models of protoplanetary discs. The corotation torque has recently received detailed attention, because of its ability to slow down, stall, or reverse type I migration. In laminar viscous disc models, the long-term evolution of the corotation torque is intimately related to viscous and thermal diffusion processes in the planet's horseshoe region. It is unclear how the corotation torque behaves in turbulent discs, and whether its amplitude is correctly predicted by viscous disc models.} {This paper is aimed at examining the properties of the corotation torque in discs where magnetohydrodynamic (MHD) turbulence develops as a result of the magnetorotational instability (MRI), considering a weak initial toroidal magnetic field.} {We present results of 3D MHD simulations carried out with two different codes. Non-ideal MHD effects and the disc's vertical stratification are neglected, and locally isothermal disc models are considered. The running time-averaged tidal torque exerted by the disc on a fixed planet is evaluated in three different disc models.} {We first present simulation results with an inner disc cavity (planet trap). As in viscous disc models, the planet is found to experience a positive running time-averaged torque over several hundred orbits, which highlights the existence of an unsaturated corotation torque maintained in the long term in MHD turbulent discs. Two disc models with initial power-law density and temperature profiles are also adopted, in which the time-averaged tidal torque is found to be in decent agreement with its counterpart in laminar viscous disc models with similar viscosity alpha parameter at the planet location. Detailed analysis of the averaged torque density distributions indicates that the differential Lindblad torque takes very similar values in MHD turbulent and laminar viscous discs, and there exists an unsaturated corotation torque in MHD turbulent discs. This analysis also reveals the existence of an additional corotation torque in weakly magnetized discs.} {Our results of 3D MHD simulations demonstrate the existence of horseshoe dynamics and an unsaturated corotation torque in weakly magnetized discs with fully developed MHD turbulence.} ",
        "introduction": "From the early stages of their formation, planets tidally interact with their natal protoplanetary disc. The tidal torque exerted by the disc on a planet drives the planet's radial migration. Several migration regimes are usually distinguished, depending on the planet's ability to clear a gap around its orbit. This ability varies with the planet-to-primary mass ratio, the disc aspect ratio (that is, the disc pressure scale height-to-radius ratio), and the disc's turbulent viscosity \\citep[for more details, see][]{crida06}. The migration of planets that do not open a gap is referred to as type I migration. It typically applies to low-mass planets, up to a few Earth masses if the central object has a solar mass. In massive protoplanetary discs, intermediate-mass planets building up a partial gap (e.g., Saturn-like planets) may undergo rapid runaway migration \\citep{mp03}. More massive planets open a clean gap and are subject to type II migration \\citep{lp86, cm07}.  The reproduction of the mass-period diagram of known exoplanets by models of planet population synthesis is very sensitive to the expression for the tidal torque driving type I migration \\citep[e.g.,][]{IdaLin4, IdaLin5, sli09, Mordasini09b}. The latter has been intensively studied in two-dimensional (2D) non-magnetized, laminar viscous disc models. In such models, the torque comprises the differential Lindblad torque and the corotation torque. The differential Lindblad torque corresponds to the angular momentum carried away by the spiral density waves the planet generates in the disc, where the flow relative to the planet becomes approximately supersonic. In discs with decreasing temperature profiles, it is a negative quantity \\citep{tanaka2002, pbck10}, which, by itself, would drive type I migration on timescales shorter than a few $\\times 10^5$ yrs \\citep[e.g.,][]{w97}. Local variations in the disc's temperature and/or density profiles, due for example to opacity transitions \\citep{mg2004} or to dust heating \\citep{hp10} may however change the sign and magnitude of the differential Lindblad torque.  The corotation torque accounts for the exchange of angular momentum between the planet and the coorbital flow. It has two components. One scales with the gradient of disc vortensity (vorticity to surface density ratio, projected along the vertical direction) across the horseshoe region \\citep{wlpi91, masset01}. The other features the entropy gradient across the horseshoe region \\citep{bm08a, pp08, mc10, pbck10}. The magnitude of the corotation torque depends on the timescales for viscous and thermal diffusion inside the horseshoe region, required to maintain the local gradients of vortensity and entropy. In the absence of diffusion processes, both gradients would be progressively flattened out by advection inside the horseshoe region, and the corotation torque would ultimately vanish. This is called saturation of the corotation torque. It may be avoided by requiring that the viscous and thermal diffusion timescales across the horseshoe region be shorter than the horseshoe libration timescale \\citep{pbk11,mc10}. When both diffusion timescales exceed the horseshoe U-turn timescale (a fraction of the libration timescale), the corotation torque is a non-linear process also known as horseshoe drag. When they become comparable to, or shorter than the U-turn timescale, the corotation torque decreases towards its value predicted by linear theory \\citep{pp09a}, referred to as the linear corotation torque.  Depending on the disc gradients and diffusion processes at the planet orbit, the amplitude of the corotation torque can be comparable to, or larger than that of the differential Lindblad torque. The impact of the corotation torque on type I migration can therefore be of considerable importance. A lot of efforts have been recently put forward to establish simple and accurate expressions for the corotation torque as well as for the differential Lindblad torque, which can be used in population synthesis models \\citep{mc10,pbk11}. These expressions were obtained for 2D non-magnetic viscous disc models.  The tidal torque has also been investigated in magnetized non-turbulent discs. The case of a 2D disc with a toroidal magnetic field has been studied by \\cite{Terquem03} through a linear analysis. She found that the differential Lindblad torque is reduced with respect to non-magnetized discs (waves propagate outside the Lindblad resonances at the magneto-sonic speed rather than the sound speed), and that there is no linear corotation torque. Instead, angular momentum is taken away from the planet by slow MHD waves propagating in a narrow annulus near magnetic resonances. These results were confirmed by \\cite{Fromangetal05} with non-linear 2D MHD simulations in the regime of strong magnetic field (the plasma $\\beta$-parameter was taken equal to 2 in their study, where $\\beta$ is the ratio of the thermal to magnetic pressures).  More recently, 2D and 3D disc models with a poloidal magnetic field were investigated by \\cite{Muto08} in the shearing sheet approximation.  It is unclear how turbulence affects the tidal torque. So far, only a couple of studies have investigated this issue. \\cite{np2004} performed 3D simulations of locally isothermal discs fully invaded by MHD turbulence due to the MRI. They found that the running time averaged tidal torque on a fixed protoplanet experiences rather large-amplitude oscillations, and its final value quite substantially differs from the torque value expected in laminar discs. Similar results were obtained by \\cite{nelson05}, who allowed the planet orbit to evolve.  A primary reason for the observed difference between the laminar torque and the time-averaged turbulent torque is that the 3D MHD simulations were not converged in time. \\cite{bl10} considered 2D isothermal discs subject to stochastic forcing, using the turbulence model originally developed by \\cite{lsa04}. They showed that, when time-averaged over a sufficient long time period, both the differential Lindblad torque and the corotation torque behave very similarly as in laminar viscous disc models. \\cite{Uribe11} have recently performed 3D MHD simulations of planet migration in weakly magnetized turbulent discs, including vertical stratification and a range of planet masses. For type I-migrating planets, they obtained a positive running time-averaged torque in disc models where the density near the planet's orbital radius becomes a slightly increasing function of radius due to the disc structuring by turbulence, suggesting the existence of an unsaturated corotation torque in such disc models.  In this paper, we revisit the properties of the tidal torque experienced by a fixed planet in disc models with hydromagnetic turbulence generated by the non-linear development of the MRI. Locally isothermal discs without vertical stratification are considered, and non-ideal MHD effects are neglected. 3D MHD simulations using two different codes have been carried out, using initial values of the plasma $\\beta-$parameter larger than or equal to 50. The physical model and numerical setup used in our simulations are described in Sect.~\\ref{sec:model}. In Sect.~\\ref{sec:cavity} we present results of simulations with an inner disc cavity. In viscous disc models, an increasing density profile yields a large positive corotation torque, which may stall type I migration. This is often referred to as a planet trap \\citep[e.g.,][]{masset06a}. Our results of turbulent simulations show a positive time-averaged tidal torque experienced by a planet fixed between the two edges of a cavity, suggesting the existence of an unsaturated corotation torque in MHD-turbulent discs. This existence is confirmed in disc models with power-law density profiles, which are described in Sect.~\\ref{sec:power}. The time-averaged tidal torque obtained in such discs is in decent agreement with the torque value predicted in 2D non-magnetized viscous disc models. The discussion in Sect.~\\ref{sec:discu} gives more insight into the existence of horseshoe dynamics in weakly magnetized turbulent discs, as well as a detailed comparison between the torque density distributions obtained in MHD-turbulent and laminar disc models. This comparison indicates the existence of an additional corotation torque in discs with a weak toroidal magnetic field, which will be investigated in a forthcoming study. Conclusions and future directions are drawn in Sect.~\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu} In this paper we have investigated the tidal torque between a planet on a fixed circular orbit and its nascent protoplanetary disc, assuming that the latter sustains turbulence driven by the non-linear development of the MRI. Our main aim has been to investigate the properties of the corotation torque in such turbulent discs.  In this initial study, we adopted simple 3D magnetized disc models with a locally isothermal equation of state, and neglected vertical stratification and non-ideal MHD effects.  We performed simulations using two different MHD codes (NIRVANA and RAMSES) over time spans of several hundred planet orbits.  In one set of simulations, we considered the torque experienced by a planet orbiting at the outer edge of an inner disc cavity, in a region with a strong, positive density gradient.  The running time-averaged torque in this case was found to be strongly positive over several hundred planet orbits, in good agreement with the results of laminar viscous disc simulations whose viscous stresses were similar to that generated by the MHD turbulence. This positive torque strongly suggests the existence of an unsaturated corotation torque in weakly magnetized turbulent disc models. It implies that the presence of an inner cavity can stall type I migration, such that a planet trap can operate in turbulent discs.  We also performed simulations of two disc models with power-law initial density profiles. We found that the running time-averaged torque of the turbulent disc models reached a stationary value after 100 to 200 planet orbits, with decent agreement between it and a laminar viscous disc model with similar viscous alpha parameter at the planet location. The relative difference between the averaged torques in turbulent and laminar disc simulations ranges from $20\\%$ to $60\\%$, depending on the code and the disc model. We find remarkably good overall agreement between the results of the two codes, and given the scales of interest (the half-width of the planet's horseshoe region is a fraction of the pressure scale height), this agreement highlights the robustness of our results.  Analysis of the velocity fields in our simulations shows the existence of horseshoe dynamics in turbulent discs with a weak toroidal magnetic field. The shape and width of the time-averaged horseshoe region much resemble those of laminar viscous discs with similar viscosity at the planet's location. Furthermore, close inspection of the averaged torque density distributions demonstrates that the differential Lindblad torque takes very similar values in turbulent and laminar viscous disc models, and there exists an unsaturated corotation torque in weakly magnetized turbulent discs. This analysis, however, also indicates the existence of an additional torque with respect to the non-magnetic case, which originates from the planet's horseshoe region. We argue that this additional torque is likely to be an additional corotation torque, possibly arising from the effects of magnetic tension.  Although we have demonstrated conclusively the existence of unsaturated corotation torques in weakly magnetized turbulent discs, our study contains a number of simplifications that need to be addressed. We have considered disc and planet parameters such that the horseshoe U-turn drift rate exceeds the amplitude of the turbulent velocity fluctuations. In the opposite case, the existence of horseshoe dynamics and a corotation torque is unknown. We have neglected vertical stratification and non-ideal MHD effects such as Ohmic resistivity, and it is fully expected that inclusion of these effects will modify the results presented here. We have also kept the planet on a fixed circular orbit and simply measured the tidal torque as a time series. In reality the planet will migrate and experience evolution in its eccentricity and inclination as a result of interaction with the turbulent disc. We will address these issues in a series of forthcoming publications, along with an in-depth investigation of the additional corotation torque in weakly magnetized discs that our simulations have uncovered."
    },
    "1107/1107.5362_arXiv.txt": {
        "abstract": "{We calculate non-Gaussianities in the bispectrum and trispectrum arising from the cubic term in the local expansion of the scalar curvature perturbation.  We compute to three-loop order and for general momenta. A procedure for evaluating the leading behavior of the resulting loop-integrals is developed and discussed. Finally, we survey unique non-linear signals which could arise from the cubic term in the squeezed limit. In particular, it is shown that loop corrections can cause $f_{NL}^{sq.}$ to change sign as the momentum scale is varied. There also exists a momentum limit where $\\tau_{NL} <0$ can be realized.}  \\arxivnumber{1107.5362} ",
        "introduction": "A variety of new experiments are poised to probe the non-Gaussianity of primordial density fluctuations with unprecedented accuracy.  These experiments include the Planck satellite~\\cite{Planck} alongside a variety of probes of large scale structure formation~\\cite{LSS}. As a result non-Gaussianity can be probed at a wide variety of scales.  It thus becomes important to determine which types of models can yield non-Gaussian fluctuations which can be probed at these experiments, and conversely, how measurements from these experiments can be used to constrain inflationary models.  Along this line, considerable work has focused on the local ansatz~\\cite{localansatz} for non-Gaussian fluctuations, which is realized in many inflationary models.  In particular it was shown that non-Gaussianity of the trispectrum is related to non-Gaussianity of the bispectrum through the inequality \\bea \\tau_{NL} \\geq \\left({6 \\over 5} f_{NL} \\right)^2 \\eea at tree-level~\\cite{Suyama:2007bg}, provided the curvature scalar is expanded to quadratic order in Gaussian fields. This became an important constraint, because it implied that measurements at upcoming experiments could potentially rule out the applicability of the local ansatz.  But subsequently it was shown in~\\cite{Sugiyama:2011jt} that this relation is modified to \\bea \\tau_{NL} \\geq {1\\over 2} \\left({6 \\over 5} f_{NL} \\right)^2 \\eea if one expands to quartic order and includes one-loop corrections.  This raises the important question of whether there is indeed a rigid constraint on the trispectrum in the local ansatz, or whether any such constraint can be violated if one computes to sufficiently high loop order.  This is especially interesting because it has been shown that in reasonable models loop contributions can dominate over tree-level contributions~\\cite{Kumar:2009ge}.  In this work, we consider a local ansatz for non-Gaussianity in a multi-field model of inflation (as non-Gaussianity is expected to be very small in single-field models with a standard kinetic term~\\cite{Maldacena:2002vr}), expanded to cubic order. We calculate the power spectrum, bispectrum and trispectrum up to three-loop order (no higher loop diagram exists for $n$-point functions with $n \\leq 4$ at cubic order in the expansion).  We will find a variety of interesting features which arise from the non-trivial scale-dependence of loop corrections (see also~\\cite{scaledep}). In particular, we find that $f_{NL}$ can change sign with momentum scale in the squeezed limit.  Moreover, $\\tau_{NL}$ can be negative in a particular limit of the external momenta.  In section 2, we review the local ansatz and the various local momentum shapes which are generated at tree-level.  In section 3, we describe the formalism for calculating loop diagrams.  In section 4, we present results for the computation of the power spectrum, bispectrum and trispectrum to cubic order (detailed calculations are presented in the appendices).  We conclude with a discussion of our results in section 5. ",
        "conclusions": "In this work, we have considered the local ansatz for non-Gaussianity in curvature perturbations. We have introduced a simple general formalism for computing higher-order loop corrections to general $n$-point correlators.  As expected, the general leading behavior of loop diagrams is the same as for tree-level diagrams, up to logarithmic corrections.  We have illustrated this procedure by computing the curvature power spectrum, bispectrum and trispectrum (expanded to cubic order in fundamental Gaussian fields) to 3-loop order.  This calculation has been made for general choices of the external momentum insertions. In particular we conclude that the relation $\\tau_{NL} \\geq (1/2)[(6/5) f_{NL}]^2$ is not necessarily obeyed for all momentum shapes.  Furthermore, there exists a squeezed limit (where the $\\tau_{NL}$ momentum structure is dominant) where $\\tau_{NL}$ can be negative if the loop-contribution dominates.  It has also been found that in the squeezed limit the sign of $f_{NL}^{sq.}$ can change as a function of scale.  This can have an interesting effect on galaxy bias.  Note that the relationship between $f_{NL}$ and $\\tau_{NL}$ depends upon the relative sizes of the external momenta. The constraint $\\tau_{NL} \\geq [(6/5) f_{NL}^{sq.}]^2$ is obtained in the elongated equilateral limit if the external momenta of the trispectrum are taken to all be of the same magnitude, while $\\tau_{NL} <0$ can be obtained if there is even a modest hierarchy in the external momenta.  The relation $\\tau_{NL} \\geq (1/2)[(6/5) f_{NL}]^2$ is only obeyed at one-loop order if the external momenta are taken to be of the same size.  It is important to remember that even if non-Gaussianities are in fact generated by the exact local ansatz (assuming the $A_i$ are constant), the bispectrum and trispectrum will not have an exactly local shape; each term in the local shape will be scaled by logarithmic corrections which depend on the particular choice of external momenta.  Unless the loop contributions are negligible, the determination of parameters such as $f_{NL}$, $\\tau_{NL}$ and $g_{NL}$ from the data will depend on what estimators are used and to which regions of momentum space they are sensitive~\\cite{Kamionkowski:2010me}. It would be interesting to determine what values of these non-Gaussian parameters would actually be yielded by standard estimators as a function of the coefficients $A_i$ of the local ansatz. \\\\ {\\bf Acknowledgements} \\\\ We thank E.~Komatsu, L.~Leblond, A.~Rajaraman and S.~Shandera for useful discussions. Preliminary results of this work have been presented at the Aspects Of Inflation Workshop and at the Cosmological non-Gaussianity: Observation Confronts Theory Workshop.  We are grateful to the organizers of these workshops, Texas A\\&M University (MIFP) and the University of Michigan (MCTP) for their hospitality. The work of JB and JK was supported in part by the Department of Energy under Grant~DE-FG02-04ER41291.  \\pagebreak \\appendix"
    },
    "1107/1107.2890_arXiv.txt": {
        "abstract": "The CHARA Array is a six 1-m telescope optical and near infrared interferometer located at the Mount Wilson Observatory in southern California and operated by the Center for High Angular Resolution Astronomy of Georgia State University. The CHARA Array has been in regular scientific operation since 2005 and now has over 55 publications in the refereed literature, including two in Science and one in Nature. The Array now supports seven beam combiners ranging from 0.5 microns up to 2.3 microns and combing from 2 to 4 beams at a time. An upgrade to a full 6 beam combiner is now underway and fringes with all six telescopes were achieved soon after the meeting. We present some of the more recent results from the CHARA-Array. ",
        "introduction": "The CHARA Array \\citep{CHARA05} has been in regular scientific operation since 2005. Many recently published results are to be found elsewhere in these proceedings \\citep{Aufdenberg2011, Baron2011, Botajian2011, Kloppenborg2011, Monnier2011, Parks2011, Rajagopal2011, Gail2011, Simon2011, Stencel2011, Touhami2011, White2011, Zhao2011}, and so in this paper we will focus on work in progress that had not been published at the time of this meeting. ",
        "conclusions": ""
    },
    "1107/1107.2318_arXiv.txt": {
        "abstract": "{In recent years new observations of pre-main sequence stars (pre-MS) with $\\mathrm{Z}\\le \\mathrm{Z}_\\odot$ have been made available. To take full advantage of the continuously growing amount of data of pre-MS stars in different environments, we need to develop updated pre-MS models for a wide range of metallicity to assign reliable ages and masses to the observed stars.} {We present updated evolutionary pre-MS models and isochrones for a fine grid of mass, age, metallicity, and helium values.} {We use a standard and well-tested stellar evolutionary code (i.e. \\texttt{FRANEC}), that adopts outer boundary conditions from detailed and realistic atmosphere models.  In this code, we incorporate additional improvements to the physical inputs related to the equation of state and the low temperature radiative opacities essential to computing low-mass stellar models.} {We make available via internet a large database of pre-MS tracks and isochrones for a wide range of chemical compositions ($\\mathrm{Z}=0.0002-0.03$), masses (M~=~$0.2-7.0$~M$_\\odot$), and ages ($1-100$ Myr) for a solar-calibrated mixing length parameter $\\alpha$ (i.e. 1.68). For each chemical composition, additional models were computed with two different mixing length values, namely $\\alpha=1.2$ and 1.9. Moreover, for $\\mathrm{Z}\\ge 0.008$, we also provided models with two different initial deuterium abundances. The characteristics of the models have been discussed in detail and compared with other work in the literature. The main uncertainties affecting theoretical predictions have been critically discussed. Comparisons with selected data indicate that there is close agreement between theory and observation.} {} ",
        "introduction": "We present a new set of pre-main sequence (pre-MS) stellar tracks and isochrones based on both the state-of-the-art input physics and outer boundary conditions. These models cover a large range of metallicities, helium abundances, masses, and ages, providing useful tools for interpreting observational data. Pre-MS tracks and isochrones collectively represent the theoretical tool needed to infer the star formation history and the initial mass function of young stellar systems. Thus, the availability of a large database with a fine grid of chemical compositions, masses, and ages will allow us to improve our understanding of the star formation process taking full advantage of the growing amount of data of very young clusters and associations in the Milky Way \\citep[e.g.][]{stolte05,delgado07,brandner08}, the Small Magellanic Cloud \\citep[e.g.][]{gouliermis06a,nota06,carlson07,gouliermis07b,sabbi07,gouliermis08,Cignoni09,cignoni10}, and the Large Magellanic Cloud \\citep[e.g.][]{romaniello04,gouliermis06b,romaniello06,gouliermis07a,dario09}.  As an example of an intriguing application of pre-MS isochrones, \\cite{Cignoni09} inferred the star formation history of NGC602, a very young stellar system in the Small Magellanic Cloud \\citep[see also][]{nota06,carlson07,gouliermis07b,sabbi07}.  From the theoretical point of view, as early understood by \\citet{hayashi61} and \\citet{hayashi63}, the evolution of hydrostatic pre-MS stars is in principle less complex than more advanced evolutionary phases from the numerical point of view, since no special algorithms are required to follow what it is essentially a quasi-static contraction. However, the pre-MS track location in the HR diagram strongly depends on the physical ingredients used in the evolutionary codes, such as the equation of state (EOS), the low-temperature radiative opacity, the outer boundary conditions, and the adopted convection treatment \\citep[see e.g.,][]{dantona93,dantona94,DM97,BCAH98,SD00,baraffe02,montalban04b}. The uncertainties due to these quantities, progressively increase as the stellar mass decreases, especially for very low-mass stars (i.e. M~$\\la 0.5$~M$_\\odot$). Therefore, to compute models as accurately as possible, it is mandatory to include the most recent updates of the above quoted physical inputs. In the past two decades, several generations of theoretical pre-MS models have been developed following improvements to the physical inputs to provide progressively more accurate theoretical predictions \\citep[e.g.][]{BCAH98,DM97,SD00,dotter07,dicriscienzo09}.  The theoretical tracks and isochrones for pre-MS stars described in the present paper have been computed with an updated version of the \\texttt{FRANEC} evolutionary code \\citep[][]{deglinnocenti08,valle09} relying on the most recent physical inputs. The models cover a wide range of metallicity ($\\mathrm{Z}~=~0.0002-0.03$), masses (M $= 0.2 -7.0$M$_\\odot$) and ages (t~=~$1-100$ Myr). The corresponding database is available on the web\\footnote{Models and isochrones are available in the theoretical plane and will soon be available in different photometric systems.}.  The input physics adopted in the calculations are described in Sect. 2, followed by a discussion of the main sources of uncertainty in the evolution of young stars. In Sects. 3 and 4, our tracks and isochrones are compared with recent evolutionary models and data. Finally, in our last section the database is described. ",
        "conclusions": "A growing amount of data of young stellar systems has prompted renewed interest in modelling the initial phase of stellar evolution.  The present paper describes a new set of pre-MS tracks and isochrones, which relies on state-of-the-art input physics (EOS, radiative and conductive opacity, atmospheric models and nuclear cross-section). To provide the astronomical community with a useful and versatile theoretical tool for the interpretation of observational data, the models have been computed for a very large and fine grid of chemical compositions. We have evaluated model data for 19 metallicities, ranging from $\\mathrm{Z}=0.0002$ to 0.03, and three different helium abundances for each Z. For a fixed chemical composition, we have made available 43 evolutionary tracks computed with three different values of the mixing-length parameter $\\alpha=1.2, 1.68$, and 1.9, in the mass range $0.2-7$~M$_\\odot$, and 36 isochrones, in the range $1-100$ Myr. For $\\mathrm{Z}\\ge 0.008$, two different initial abundances of deuterium have been adopted. The database is available on the web\\footnote{ \\url{http://astro.df.unipi.it/stellar-models/}}. Models are compared in detail with other computations available in the literature, while the comparison with selected observational data shows good agreement."
    },
    "1107/1107.2404_arXiv.txt": {
        "abstract": "We present an analysis of the $B$-band and $V$-band rise-time distributions of nearby Type Ia supernovae (SNe~Ia). Drawing mostly from the recently published Lick Observatory Supernova Search sample of SNe~Ia \\citep{ganeshalingam10}, together with other published nearby SNe~Ia with data starting at least one week before maximum light, we use a two-stretch template-fitting method to measure the rise and decline of $BV$ light curves.  Our analysis of \\noo\\ SNe with high-quality light curves indicates that the longer the time between explosion and maximum light (i.e., the rise time),  the slower the decline of the light curve after maximum. However, SNe with slower post-maximum decline rates have a faster rise than would be expected from a single-parameter family of light curves, indicating that SN~Ia light curves are not a single-parameter family of varying widths. Comparison of the $B$-band rise-time distribution for spectroscopically normal SNe~Ia to those exhibiting high-velocity spectral features indicates that high-velocity (HV) SNe~Ia have shorter $B$-band rise times compared to their spectroscopically normal counterparts.  After normalising the $B$-band  light curves to \\dmb\\ = 1.1~mag (i.e., correcting the post-maximum decline to have the same shape as our template), we find that spectroscopically normal SNe~Ia have a rise time of $18.03 \\pm 0.24~{\\rm d}$, while HV SNe have a faster $B$-band rise time of $16.63 \\pm 0.29~{\\rm d}$. Despite differences in the $B$ band, we find that HV and normal SNe~Ia have similar rise times in the $V$ band. Furthermore, the initial rise of a SN~Ia $B$-band light curve follows a power law with index  $2.20^{+0.27}_{-0.19}$,  consistent with a parabolic rise in flux predicted by an expanding fireball toy model.  We compare our early-time $B$-band data to models for the predicted signature of companion interaction arising from the single-degenerate progenitor scenario. There is a substantial degree of degeneracy between the adopted power-law index of the SN light-curve template, the rise time, and the amount of shock emission required to match the data. ",
        "introduction": "Type Ia supernovae (SNe~Ia) are believed to be the result of a thermonuclear runaway explosion of a C/O white dwarf (WD) approaching the Chandrasekhar limit  (see \\citealt{hillebrandt00} for a review). Explosive nucleosynthesis up to \\nic releases $\\sim 10^{51}{\\rm ~erg}$, unbinding the progenitor. The subsequent light curve is powered by injection of energy from the radioactive decay of \\nico; the $\\gamma$ rays degrade to longer wavelengths as they diffuse out through the expanding ejecta.  The well-established relationship between the light-curve width and the luminosity at peak brightness allows SNe~Ia  to be ``standardizable candles\" at optical wavelengths \\citep{phillips93}, and possibly almost standard candles in the infrared \\citep{krisciunas04, wood-vasey08, folatelli10}. Application of the width-luminosity relation along with colour information to SNe~Ia over cosmological scales has led to the discovery of the accelerating expansion of the Universe (\\citealt{riess98, perlmutter99}; see also \\citealt{astier06, riess07, wood-vasey07, kowalski08, hicken09b,amunullah10}), indicating either the presence of ``dark energy'' having a negative pressure or a failure of general relativity on the largest scales. Recent work has also shown that including spectral flux ratios may also aid in reducing Hubble-diagram residuals \\citep{bailey09, blondin11}.  Despite the successful cosmological application of SNe~Ia, the SN community lacks a clear understanding of the nature of their progenitor systems. Possible scenarios include a single-degenerate WD paired with a red-giant post-main-sequence star undergoing stable mass transfer until $M_{\\rm WD}$ approaches $M_{\\rm Ch} \\approx 1.4$~\\Msun \\citep{whelan73,livio00}, a double-degenerate WD merger that reaches or surpasses $M_{\\rm Ch}$ \\citep{webbink84, iben94}, or the result of a sub-Chandrasekhar explosion of a WD steadily burning helium accreted from a companion \\citep{shen09}. Each scenario carries with it tentative observational evidence and hardships imposed by unrealised theoretical predictions.  While considerable attention has been focused on the post-maximum decline of the SN~Ia light curve, less has been paid to the rise of the SN~Ia from explosion to maximum brightness. This is due to the dearth of data in the days following the SN explosion and the intrinsic difficulty of finding SNe~Ia shortly after explosion. The Lick Observatory Supernova Search (LOSS) has been successful at finding SNe in the nearby Universe (redshift $z < 0.05$) and the Sloan Digital Sky Survey (SDSS) Supernova Survey \\citep{frieman08} has found and monitored 391 SNe at moderate redshift ($0.03 \\le z \\le 0.35$). Higher redshift searches such as the SuperNova Legacy Survey \\citep{astier06} and ESSENCE \\citep{wood-vasey07} benefit from cosmological time dilation, allowing the SN rise to be spread out over more days in the observer's frame.  A first attempt at quantifying SN~Ia rise times by \\cite{pskovskii84} analysed 54 historical light curves gathered from photographic plates and visual $B$ magnitude estimates, finding a typical rise of 19--20~d. Recent attempts have made use of more reliable data taken with charge-coupled devices (CCDs) which are linear in translating photons to counts over a large dynamic range. A seminal paper by \\citet[][hereafter R99]{riess99b} used observations of nearby SNe~Ia to find a $B$-band rise time of $19.5 \\pm 0.2~{\\rm d}$ for a normal SN~Ia having a decline of 1.1 mag between maximum light and 15~d after maximum in the $B$ band (i.e., $\\dmb =1.1$ mag; Phillips 1993). Studies using data on higher redshift SNe~Ia have found a rise time in concordance with that of R99 \\citep{aldering00,conley06}, advancing the notion that there is limited, if any, evolution in SN~Ia properties from high to low redshifts. Most recently, \\cite[][hereafter C06]{conley06} found a rise of $19.10^{+0.18}_{-0.17} {\\rm~(stat)} \\pm 0.2 \\rm{(sys)}$~d using SNLS data.  In an analysis of $B$- and $V$-band photometry of eight nearby SNe~Ia with excellent early-time data, \\citet[][hereafter S07]{strovink07} found tentative evidence for two populations of $B$-band rise times. After correction for light-curve decline rate, S07 found that three SNe rise in $18.81 \\pm 0.36$ d and five SNe rise in $16.64 \\pm 0.21$ d. More recently, \\citet[][hereafter H10]{hayden10a} analysed $B$- and $V$-band photometry 105 SNe~Ia from the SDSS sample and found that SNe~Ia come from a rather broad distribution of $B$-band rise times, with evidence indicating that slower declining events (e.g., more luminous SNe) have some of the fastest rise times. This result has significant implications for light-curve fitting techniques that rely on a single-parameter family of light curves and theoretical modeling of SNe~Ia. {H10 find an average $B$-band rise time of $17.38 \\pm 0.17$~d, a departure from the results of R99 and C06 which H10 trace back to differences in fitting methods.  In this paper, we analyse available data on nearby SNe~Ia to measure the rise times, relying heavily on the recently released LOSS sample \\citep{ganeshalingam10}. Previous analyses such as those of S07 and H10 use combined results from both $B$- and $V$-band photometry to measure the $B$-band rise time. In this paper, we compare the $B$-band rise time measured in the $B$ and $V$ bands and find that such a combination may not be appropriate. Instead, we present independent analysis of the $B$ and $V$ bands. Specifically, we define the rise time in a photometric band as the elapsed time between explosion and maximum light for that particular band. Nearby SNe offer the advantage of being able to be monitored by small-aperture telescopes and benefit from not requiring significant $K$-corrections, which at early times are ill defined because of a lack of available spectra to model the SN spectral energy distribution (SED). ",
        "conclusions": "We have presented an analysis of the rise-time distribution of nearby SNe~Ia in the $B$ and $V$ bands. Using a two-stretch fitting technique, we find that the SN rise time is correlated with the decline rate in the sense that SNe with broader light curves post-maximum (i.e., light curves with small \\dmb /\\dmv) have longer rise times.  While SN 1991bg-like objects could not be fit well by our two-stretch fitting procedure, we found that restricting our analysis to the pre-maximum data for the SN 1991bg-like SN 1999by gives a rise time of $13.33 \\pm 0.40$~d. This follows the expected trend of a fast rise leading to a fast decline.  Using a sample of 105 SDSS SNe at intermediate redshifts ($0.037 \\leq z \\leq 0.230$), H10 find that there is a great diversity in $B$-band rise times for a fixed \\dmb, and that the slowest declining SNe tend to have the fastest rise times. While we do see evidence for  scatter in rise times with \\dmb, we find a strong correlation between \\dmb\\ and rise time in the sense that slowly declining SNe~Ia have longer rise times. This correlation is strong in the $B$ band, but weaker in the $V$ band.  The discrepancy with  H10 may be a result of combining $B$- and $V$-band stretches which was avoided in this analysis.  We find that a single value of the stretch does not adequately describe the rising and falling portions of our light curves. Similar to the results in H10, more luminous SNe with longer fall times have shorter rise times than one would expect from a single-stretch prescription for light curves. This is especially evident in our analysis of $V$-band light curves. However, we reiterate that while luminous SNe have shorter rise times than expected from a single-stretch prescription, they still have longer rise times than less luminous SNe, contrary to the findings of H10.  R99 and C06 found a fiducial $B$-band rise time of $\\sim 19.5$~d for a ``typical\" SN~Ia. In a departure from previous results, H10 report an average rise time of $17.38 \\pm 0.18$~d. We find a fall-stretch corrected (i.e., corrected to have a post-maximum fall of \\dmb\\ = 1.1 mag) $B$-band rise time of $18.03 \\pm 0.24$~d for spectroscopically normal SNe and $16.63 \\pm 0.30$~d for HV SNe. Our $B$-band rise time for spectroscopically normal SNe with \\dmb\\ = 1.1 mag is in agreement with the average rise time found by H10 at the $2.2 \\sigma$ level. When correcting our $V$-band light curves to a post-maximum fall of \\dmv\\ = 0.66 mag, we find a fall-stretch corrected $V$-band rise time of $20.23 \\pm 0.44$~d.  After correcting for post-maximum decline rate, HV SNe~Ia have faster rise times than normal SNe~Ia in the $B$ band, but similar rise times in $V$. We find a $\\sim 3\\sigma$ difference in the fall-stretch corrected rise time between HV and normal SNe in the $B$ band.  The rise minus fall (RMF) distributions (not corrected for stretch) of the two subclassifications show significant differences in the two populations. The peak values of the distributions are offset by $1.22 \\pm 0.34$~d in the $B$ band, with HV objects having a faster rise time. A K-S test indicates a $\\sim 0.01$\\% probability that HV SNe come from the same parent RMF population as normal SNe in the $B$ band. Despite differences in the $B$ band, we see no evidence of a difference in the RMF populations in the $V$ band.  Based on the model presented by \\cite{foley11} and the models of KP07, we offer a possible qualitative explanation for why HV SNe should have a different $B$-band rise-time distribution than normal SNe, but similar $V$-band distributions. The physical origin of the difference is possibly rooted in the different opacity mechanisms at work in the $B$ and $V$ bands. Line blanketing from Fe-group elements is the dominant source of opacity for wavelengths shorter than $\\sim 4300$ \\AA, the peak of the $B$ band. At longer wavelengths, such as the $V$ band, electron scattering is the dominant source of opacity. Rapidly moving ejecta, as is the case with HV objects, will broaden absorption features, diminishing the $B$-band flux without affecting the $V$-band flux. Models from KP07 show that all other things being equal (e.g., Ni mass, kinetic energy), this leads to faster light curves with larger \\dmb\\ for HV objects. However, the enhanced opacity at short wavelengths also affects \\dmb, complicating the application of the models to our result. Further modeling and observations of HV SNe~Ia are required to shed light on the photometric differences of this spectroscopic subclass.  We fit the earliest data in our sample ($\\tau \\leq -10~\\rm{d}$) to find that the flux rises as a power law with index $2.20^{+0.27}_{-0.19}$. This is consistent with the unimpeded, free expansion of the expanding fireball toy model that predicts an index of 2. However, a preliminary analysis of SN~2009ig in the range $-15 < \\tau < -9$~d shows evidence of significant colour evolution, contrary to the assumption of little to no colour evolution in the expanding fireball model. H10 find similar colour evolution in an analysis of the $B-V$ colour curve of SDSS SNe~Ia and derive an expected power-lax index of 4.  We compare our early-time $B$-band data as an ensemble to models \\citep{kasen10} of shock interaction produced from SN ejecta colliding with the mass-donating companion in the single-degenerate progenitor scenario. When relaxing our assumptions on the functional form of the early-time template light-curve behaviour (i.e., changing the rise time or power-law index of the rise), we find that our data require some amount of shock interaction to remove systematic trends. This indicates a level of degeneracy between the adopted template rise time, the power-law index, and the amount of shock interaction required to match the data.  Future surveys with high-cadence search strategies will provide well-sampled SN light curves starting days after explosion, substantially adding to the sample of rise-time measurements. Complemented with spectroscopic follow-up observations, analysis of the RMF distribution can be further broken into different spectroscopic classes to provide insights into the underlying populations and the physics that differentiates them."
    },
    "1107/1107.0415_arXiv.txt": {
        "abstract": "In this paper we model the evolution of PrePlanetary Nebula (PPN) and Planetary Nebula (PN) morphologies as a function of nebular age. The aim of the work is to understand if shape transitions from one evolutionary phase to the other can be driven by changes in the parameters of the mass-loss from the central star. We carry out 2.5D~hydrodynamical simulations of mass-loss at the end stages of stellar evolution for intermediate mass stars.  Changes in wind velocity, mass-loss rate and mass-loss geometry are tracked. We focus on the transition from mass-loss dominated by a short duration jet flow (driven during the PPN phase) to mass-loss driven by a spherical fast wind (produced by the central star of the PN). Our results show that while jet driven nebulae can be expected to be dominated by bipolar morphologies, systems that begin with a jet but are followed by a spherical fast wind will evolve into elliptical objects.  Systems that begin with an aspherical AGB wind evolve into butterfly shaped nebula with, or without, a jet phase. In addition, our models show that spherical nebulae are highly unlikely to derive from either bipolar PPN or elliptical PN over relevant time scales.  The morphological transitions seen in our simulations may however provide insight into the driving mechanisms of both PPN and PN as point to evolutionary changes in the central engine. ",
        "introduction": "Planetary Nebulae (PN) expand at a~few~10\\,km\\,s$^{-1}$ away from evolved intermediate-mass stars (with main sequence progenitor masses $\\sim\\,$1$-$8\\,M$_{\\odot}$). These nebulae exhibit a rich variety of morphologies with the main categories being spherical, elliptical and bipolar \\citep[for a review see][]{balick02}. % Spherical PN have been explained by time-dependent interacting stellar wind models (ISW;  Kwok, Purton, \\& Fitzgerald 1978) that describe the nebula as the collision of a fast ($\\sim\\,$1000\\,km\\,s$^{-1}$) tenuous (between about 10$^{-8}$ and 10$^{-7}$\\, M$_{\\odot}$\\,yr$^{-1}$) stellar wind driven by the hot central star of PN % (CSPN) with a dusty slow ($\\sim\\,$10\\,km\\,s$^{-1}$) dense ($\\sim\\/$10$^{-4}$\\,M$_{\\odot}$\\,yr$^{-1}$) wind, or circumstellar envelope, ejected during the earlier AGB phase. The nebulae are thus formed when the fast wind catches up with the AGB wind, driving a shocked shell through it.  The formation of bipolar PN and jets has been addressed by generalized ISW models (GISM) in which the ejected AGB envelope is characterized by either a toroidal density distribution or an aspherical velocity field ~(Kahn \\& West 1985; \\citealp{balick87,icke88,huggins07}; Frank \\& Mellema~1994; \\citealp{icke88}; Icke, Preston, \\& Balick, 1989; Mellema, Eulderink, \\& Icke, 1991; Icke, Balick, \\& Frank, 1992; \\citealp{frank94}; Dwarkadas, Chevalier, \\& Blondin, 1996). In some of these models, asymmetries in AGB envelopes develop as a result of a binary stellar system, in which an AGB star interacts with a companion \\citep{arredondo,edgar}, or an AGB star and a companion share a common envelope evolution (\\citealp{iben91,livio93,soker97}; \\citealp{sandquist98}; \\citealp[and references therein]{soker98}; \\citealp{demarco03,nordhaus07,passy12}). Note, however, that AGB envelopes generally show spherical % shapes (Bujarrabal \\& Alcolea 1991; Kahane \\& Jura 1994; \\citealp{stanek95,groeneweger96}). Some AGB observations show seeds of asymmetry.  The original GISW model (see e.g. Balick \\& Frank 2002) assumed a spherical fast wind driven from the CSPN expanding into an aspherical AGB wind. While these models were able to generate hot, low density jet flows via shock focusing at the inner shock, they were less successful at recovering dense colder jets like those seen in YSOs.  Since many PPN and young PN show evidence for such narrow jets, Sahai \\& Trauger (1998) suggested that collimated PN flows created at, or near, the central engine were the real drivers of early PN morphology. Such jet formation has been proposed as a natural consequence of interacting binary evolution (\\citealp[][2004]{lee03}; \\citealp{akashi08,lee09}). \\citet{lee03} % carried out numerical simulations and synthetic observations of a jet (or collimated fast wind) interacting with a spherical AGB wind. They found that both the dynamics and emission of PPN, and young PN, shells depend on the velocity and geometry of the jet. % Using a similar approach \\citet{akashi08} % modeled the effect of short-duration jets expanding into a spherical, slow AGB wind on the observed shape of evolved PN. The simulations of \\citet{akashi08} covered timescales about an order of magnitude longer than the ones of \\citet{lee03}. \\citet{akashi08} found that the AGB mass-loss history also has a profound influence on the nebular structure, and that the head of the lobes move sufficiently fast to excite some visible emission lines (see \\citealp{lee03} for detailed results of optical forbidden line emission which is excited at the heads of the lobes during the PPN phase). % % These studies did not address the question as to how the older, AGB spherical outflows would respond to ``fresh'', collimated outflows from the PN central engine.  In this paper we focus on a theoretical exploration of the role that changing mass-loss geometries from the central engine can play in driving changes in nebular morphology.  If PPN drive collimated outflows off the central engine but PN drive primarily spherical outflows then how will the nebula respond as it expands? The aim is to understand how nebulae change form as their hydrodynamic driving changes. Such process is of interest because it might help to understand the history of the central engine which may be a single star, a recently processed common envelope system or a binary with an accretion disk. Thus there is the hope that one might be able to use morphological clues in an individual nebula to recover something of the history of the central object in terms of its mass-loss and the conditions which govern it.  We note that there is a broader observational motivation for this study.  The classifications of PPN and PN by morphological type appear to show significant statistical differences. The general trends indicate that (i) around~70 to~90\\% of PPN and young PN are bipolar or more complex \\citep{sahai98,sahai11}, (ii) about 20\\% of mature PN are round, 70\\% are elliptical and about~10\\% are bipolar (see \\citealp{miszalski09a,parker06}; \\citealp{lagadec11}; the Planetary Nebula Image Catalog of Bruce Balick\\footnote{http://www.astro.washington.edu/users/balick/PNIC/} and references therein).  There are however considerable uncertainties in these classifications, e.g. (i) telescope resolution limitations; (ii) stellar population selection effects; (iii) differences in nebular class definitions (for example ring-like nebulae that are actually toroids included in spherical morphological classes); (iv) the well-established correlation between PN morphology and the progenitors' mass, i.e., bipolar and ring-like PN generally found at lower galactic scale heights, characteristic of higher-mass progenitors \\citep[see e.g.][]{corradi95,kastner96,manchado11}. These factors, along with the fact that most attempts to classify PN by morphological type do not consider the ages of the nebulae (but see \\citealp{sahai98,sahai11} for specific studies on young PN), complicate the interpretation of temporal changes in nebular morphology.     If the imbalance between PPN and PN statistics holds up under further scrutiny then it implies that some objects may undergo a shape transformation as they expand. In particular, one can imagine a scenario in which bipolar PPN hydrodynamically transform into elliptical or rounder PN. Such a transformation could be caused by different nebular driving mechanisms during the early and late nebular evolutionary phases. % % %  Irrespective of issues of observational statistics, in this paper we take on the issue of morphological changes in PPN/PN from a theoretical perspective. Starting from a set of basic physical ingredients and using simulations, we aim  to understand how shape transitions might occur during the evolution of an individual post-AGB system and how such morphological transitions  depend on the parameters of the mass-loss process of the central star. The morphological evolution of single sources should be studied from first principles to better inform the interpretation of observations. Thus we consider the influence of an initial collimated jet phase during the PPN and its effect on the previously deposited circumstellar AGB wind. The jet phase is followed by a classical fast wind with a spherical geometry.  Different distributions of the AGB wind's velocity and density are explored as are other parameters for the jet/wind interactions.  Our choices for initial and boundary conditions are based on a reasonable balance between observational expectations and computational expediency. We note \\citet{soker90} has studied PN ansae using a similar wind/jet sequence, but here we carry out a more extensive study and explore a range of parameters across both PPN and PN evolution.  This paper is organized as follows: in section~\\ref{model} we describe the methodology and numerical code used in the study, as well as our model and implementation of the AGB wind and post-AGB stellar outflows. The results of the simulations are presented in section~\\ref{results}. In section~\\ref{discu} we compare our results with the general trends found by observations by means of synthetic emission maps and discuss how the results bear on issues of PN shaping. Conclusions are presented in section~\\ref{conclu}. ",
        "conclusions": "\\label{conclu}  We have carried out a series of 2.5D hydro simulations designed to study how changing the mass-loss properties of PN progenitor stars affects the evolution of nebular morphology. The stellar mass-loss evolution has been followed starting at the AGB phase, evolving through the PPN phase and finishing at the mature PN phase.  Our simulations show that a young bipolar PPN transforms into an older elliptical PN when an initially spherical AGB envelope interacts with a short duration jet, driven from the central star and  followed by an isotropic fast wind.  Thus we have demonstrated how changing mass-loss geometries from the central engine can drive evolutionary changes in nebular morphology.  Our theoretical results complement GISW studies.  A fast spherical CSPN wind driving into a slow spherical AGB wind has been thought to create a spherical nebula.  We have shown that if this interaction is preceded by a collimated PPN jet phase then an elliptical PN will result. A fast spherical CSPN wind driving into a slow aspherical AGB wind has been thought to create an aspherical nebula and our results show that including an earlier PPN jet phase does not affect the eventual butterfly morphology that results.  Our theoretical result that bipolar PPN can evolve into elliptical PN is consistent with general observational trends suggesting that PPN appear to favor bipolar morphologies while mature PN appear to be more elliptical. This  result  has to be considered along side of the complementary  observational trend that   bipolar PN are found at typically lower Galactic scale heights than ellipticals,  suggesting a correlation between bipolarity and higher-mass (shorter lifetime) progenitors \\citep{corradi95,kastner96,manchado11}.  We also have found that once mature PN become elliptical they do not evolve to become spherical nebulae over the relevant PN lifetimes, so the origin of mature spherical PN (19\\% of all PN; \\citealp{parker06,miszalski08}) does not seem to follow the initially aspherical pathway studied herein.  The aspect ratio evolution of our model nebulae suggests that bipolar PN with projected aspect ratios~$\\ga\\,$4 may result from CSPN which during the post-AGB phase develop a brief jet, but no fast spherical winds, or relatively weak ones (e.g., symbiotic Mira systems like R~Aqr, Mz~3, \\citealp{Mira}, and references therein; Mz~3, \\citealp{schwarz92}).  When the initial jet drives into a toriodal AGB density distribution the result is a bipolar PPN that evolves into a mature bipolar nebula with a narrow waist. The toroidal AGB density distribution in this case dominates the dynamics and shape of any subsequent outflow phases. At late times the asymptotic geometry of this nebula is a traditional butterfly nebula when seen in projection onto the sky. We also find that the interaction of a spherical post-AGB fast wind with an AGB envelope which has an aspherical velocity field yields a homologously expanding elliptical nebula.  One of our simulations included a quiescent interlude between the jet and the fast wind episodes. We find the corresponding synthetic emission map shows small bright features which resemble observations of axial knots or FLIERS within PN. A future avenue for these studies would be to explore when and how knots form as the result of jet material collecting at the head of the flow."
    },
    "1107/1107.5006_arXiv.txt": {
        "abstract": "We employ a coarse-graining approach to analyze nonlinear cascades in Boussinesq flows using high-resolution simulation data. We derive budgets which resolve the evolution of energy and potential enstrophy simultaneously in space and in scale.  We then use numerical simulations of Boussinesq flows, with forcing in the large-scales, and fixed rotation and stable stratification along the vertical axis, to study the inter-scale flux of energy and potential enstrophy in three different regimes of stratification and rotation: (i) strong rotation and moderate stratification, (ii) moderate rotation and strong stratification, and (iii) equally strong stratification and rotation.  In all three cases, we observe constant fluxes of both global invariants, the mean energy and mean potential enstrophy, from large to small scales. The existence of constant potential enstrophy flux ranges provides the first direct empirical evidence in support of the notion of a cascade of potential enstrophy.  The persistent forward cascade of the two invariants reflects a marked departure of these flows from two-dimensional turbulence. ",
        "introduction": "The Kolmogorov \\cite{Kolmogorov41a} and Kraichnan \\cite{Kraichnan67} theories of three- and two-dimensional Navier-Stokes turbulence have served as a benchmark in the understanding of fluid turbulence and as fundamental tests for the accuracy of simulations. The Boussinesq approximation of the compressible Navier-Stokes equations in a rotating frame give a fairly accurate description of the flow dynamics over much of the Earth's oceans and atmosphere but are prohibitively expensive to simulate in detail over global scales. Guided by the success of the Kolmogorov/Kraichnan theories, it would be useful to develop a statistical phenomenology of the small scales of Boussinesq flows to gain an understanding of the physics, serve as benchmarks to test simulations, and offer parameterizations which could eventually be useful in practical modeling of geophysical flows. In addition to global energy, the inviscid Boussinesq equations conserve local potential vorticity (hereafter, PV) and global potential enstrophy, thus offering more complexity than incompressible Navier-Stokes dynamics. Charney \\cite{Charney71} addressed one limiting case of the Boussinesq approximation, namely the quasi-geostrophic limit of strong rotation and strong stratification, and showed that the conservation of both energy and the quadratic potential enstrophy in such flows constrained energy to cascade to the large scales as in 2D turbulence \\cite{Kraichnan67}.   Conservation of potential enstrophy is believed to play a fundamental role in the dynamics of the atmosphere and oceans \\cite{Vallis06}. Understanding its function in nonlinear scale interactions would appear to be essential for extending Kolmogorov's theory to Boussinesq flows with rotation and stratification. Herring et al. \\cite{Herringetal94} studied the cascade properties of potential enstrophy in turbulence simulations with a passive scalar in the absence of rotation and stratification. The authors concluded that because potential enstrophy in their simulation is not a quadratic, but a quartic invariant, the usual Kolmogorov-like arguments for a cascade and an inertial range do not apply. In fact, their potential enstrophy budget shows the direct action of viscous-diffusion terms at \\emph{all scales}, even in the limit of very small viscosity and diffusivity. Their figure 15 shows that potential enstrophy dissipation peaks at the largest scales (see also their figure 14 and the discussion on pp. 37 and 43). Thus, a pure inertial range of potential enstrophy flux is precluded because of contamination by dissipation at all scales. However for Boussinesq flows, with strong rotation and/or strong stratification, Kurien et al. \\cite{Kurienetal06} (hereafter, KSW06) derived analytically, starting from the evolution of the two-point correlation of potential vorticity, a flux law for potential enstrophy which is analogous to Kolmogorov's 4/5-law for energy flux in 3D incompressible turbulence. The so-called 2/3-law of KSW06 implies that an inertial cascade of potential enstrophy can exist in three limiting cases: (i) strong rotation with moderate stratification (ii) moderate rotation with strong stratification and (iii) strong rotation and strong stratification. By ``moderate'' we mean that the rotation (stratification) frequency is of the same order as the non-linear turbulence frequency; whereas by ``strong'' we mean that the respective frequencies are much faster than the non-linear timescale. In  the above three regimes, KSW06 showed that potential enstrophy, generally a quartic quantity, becomes quadratic which results in the localization of viscous-diffusion terms to the smallest scales, thus allowing for an inertial range of scales dominated by the flux of potential enstrophy. Furthermore, KSW06 suggested that in the absence of strong rotation and/or strong stratification, when potential enstrophy reverts to being quartic, viscous-diffusion effects may contaminate all scales, in agreement with the conclusions of \\cite{Herringetal94}.  The existence of an inertial cascade range for potential enstrophy is far from obvious and remains an unsettled issue. To date, there has been no empirical demonstration of KSW06's results on the constant flux range of potential enstrophy in Boussinesq flows. In \\cite{Kurienetal08} a phenomenology and supporting data from Boussinesq simulations with equally strong (non-dimensional) rotation and stratification were presented to show that conservation of quadratic potential enstrophy could constrain the spectral distribution of energy in the large wavenumbers. In \\cite{Kurien11} the analysis and phenomenology of \\cite{Kurienetal08} was extended to include the two other limiting cases (i) and (ii) above, to show that indeed in all parameter regimes with linear PV, the downscale flux of (quadratic) potential enstrophy can constrain the scale-distribution of energy. While the two studies \\cite{Kurienetal08,Kurien11} did not measure potential enstrophy flux, their numerical findings confirmed phenomenologically predicted scalings of the energy spectra assuming constant downscale fluxes of potential enstrophy. Motivated partially by those results, the present work uses data from \\cite{Kurien11} to directly measure the cascade of potential enstrophy for the first time.  Some previous studies have highlighted the interest in energy fluxes. The results from \\cite{Lindborg06,Brethouweretal07} of purely stratified flow computed in small-aspect-ratio domains with forcing of the horizontal velocity only, did not show convincing constant energy fluxes (see figure 18 of \\cite{Lindborg06} and figure 15 of \\cite{Brethouweretal07} which plot energy flux on a log-log scale). On the other hand \\cite{Waite11} provides evidence of scale-independent fluxes of kinetic and potential energy in purely stratified flow in unit aspect-ratio. It should be noted that these studies used forcing, stratification, and domain aspect-ratio different from ours. Furthermore, the simulations of \\cite{Lindborg06,Brethouweretal07,Waite11} were not reported to be in the linear PV regime and, therefore, our results may not apply to their flows. Most importantly, none of these previous studies computed or analyzed the potential enstrophy flux, which constitutes an essential part of our work.    In this Letter, we present a very general framework for analyzing nonlinear scale interactions in Boussinesq flows. The coarse-graining approach we utilize allows for probing the dynamics simultaneously in space and in scale.  Motivated by the work of \\cite{Kurienetal06,Kurienetal08,Kurien11}, we then measure fluxes of energy and potential enstrophy across scales from simulations in three distinct limits of rotation and stratification. Our results show constant and positive fluxes of the two quadratic invariants, indicating simultaneous persistent downscale cascades of both quantities in all three cases.  Our measurements of potential enstrophy flux are a novel contribution of this Letter and constitute the first empirical confirmation of analytical results by \\cite{Kurienetal06}. Furthermore, our evidence of a scale-independent energy flux is significant because it conveys that a cascade should persist to arbitrarily small scales at asymptotically high simulation resolutions. ",
        "conclusions": "The two main results presented in this Letter are (i) energy and potential enstrophy budgets which resolve the dynamics simultaneously in space and in scale and (ii) concurrent and persistent cascades of energy and potential energy to small scales in three extreme cases of rotating and stratified Boussinesq flow simulations. The numerical results on constant fluxes of potential enstrophy constitute the first direct empirical evidence in support of analytical results by \\cite{Kurienetal06}. Our findings show a clear departure of the flows we study from 2-dimensional turbulence and should be incorporated in any phenomenological treatments of strongly rotating and/or stratified Boussinesq flows. In a longer forthcoming work \\cite{AluieKurien_long}, we shall refine our analysis to study anisotropy of these cascades, their pointwise and scale-locality properties, and quantify contributions from vortical and wave components.     \\vspace{0.4cm} \\noindent {\\small {\\bf Acknowledgements.} This research used resources of the Argonne Leadership Computing Facility at Argonne National Laboratory, which is supported by the Office of Science of the US  DOE under contract DE- AC02-06CH11357. HA acknowledges partial support from NSF grant PHY-0903872 during a visit to the Kavli Institute for Theoretical Physics. This research was performed under the auspices of the US DOE at LANL under Contract No. DE-AC52-06NA25396. HA was supported by the LANL/LDRD program and by the DOE ASCR program in Applied Mathematical Sciences. SK received partial funding from NSF program Collaborations in the Mathematical Geosciences: NSF CMG-1025188. }"
    },
    "1107/1107.2556_arXiv.txt": {
        "abstract": "{We present deep V and I  photometry of the isolated dwarf galaxy VV124=UGC4879,  obtained from archival images taken with the \\emph{Hubble Space Telescope - Advanced Camera for Surveys}. In the color-magnitude diagrams of stars at distances larger than $40\\arcsec$ from the center of the galaxy, we clearly identify for the first time a well-populated old Horizontal Branch (HB). We show that the distribution of these stars is more extended than that of Red Clump stars. This implies that very old and metal poor populations becomes more and more dominant in the outskirts of VV124. We also identify a massive ($M=1.2\\pm 0.2\\times10^4~M_{\\sun}$) young (age$= 250\\pm 50$~Myr) star cluster (C1), as well as another of younger age (C2, $\\la 30\\pm 10$~Myr) with a mass similar to classical open clusters  ($M\\le 3.3\\pm0.5\\times10^3~M_{\\sun}$). Both clusters lie at projected distances smaller than 100~pc from the center of the galaxy.\\\\} ",
        "introduction": "\\label{intro}  The dwarf galaxy VV124=UGC4879 has been recently recognized to lie at the outer fringes of the Local Group (LG), near the turn-around radius \\citep{k08,tik}. It is extremely isolated, being at $\\simeq$1.3~Mpc from the nearest giant galaxy (the Milky Way) and at $\\simeq$0.7~Mpc from the nearest dwarf (Leo~A).  In \\citet[][Pap-I hereafter]{vv124} we presented new deep wide-field photometry of VV124, which allowed us to confirm and refine the distance estimate by \\citet{k08}, to study in some detail the stellar content of the galaxy and, in particular, its structure out to large radii, revealing the presence of low surface brightness structures in the outskirts. We also presented the first detection of \\HI ~in VV124, obtained with the  Westerbork Synthesis Radio Telescope (WSRT). $10^6~M_{\\sun}$ of \\HI ~were found, with a maximum column density of $\\ga 50\\times10^{19}$~cm$^{-2}$. It is worth noting that the distribution of \\HI ~is less extended than the stellar body of the galaxy.  \\citet[][J11 hereafter]{j11} presented Hubble Space Telescope (HST) - Advanced Camera for Surveys (ACS) observations of VV124, resolving very well the stars down to $F814W\\simeq 27$ also in the most central region of the galaxy. In particular, they studied the stellar content in the innermost $40\\arcsec$ (see their Fig.~1), deriving the star formation history (SFH) in that region by means of the synthetic color-magnitude diamgram (CMD) technique \\citep[see, for example][and references therein]{rizzi,cignoni}. J11 concluded that most of the SF ($\\simeq$ 93  percent of the total, in mass) occurred at the earliest epochs, between 14 and 10 Gyr ago. After this, the SF ceased until 1 Gyr ago, when $\\simeq$6 percent was formed between 1 and 0.5 Gyr ago and nearly 1 percent in the last 0.5 Gyr. The derived SFH is similar to that observed in other nearby dwarf irregular/transition-type galaxies \\citep[see, for example][]{weisz}. The prevalence of old stars seems to increase at larger distances from the center (see also Pap-I). From the morphology of the blue-loop region, J11 inferred that the metallicity of the youngest stars is $[M/H]=-0.7$.  In the present contribution, we use the same observational material of J11 to address three issues raised in Pap-I and not considered by J11. These are, in particular:  \\begin{itemize}  \\item The tentative detection of an extended Horizontal Branch (HB), including  stars in the color range of RR Lyrae variables and blue HB stars, in the CMD of the less-crowded outer regions of the galaxy (see Fig.~9 of Pap-I). While J11 report that a population of HB stars reaching $V-I\\sim 0.4$ is likely present in their CMD, its clean identification and quantification is hampered by the overlap with the young main sequence (MS) in the central region they consider. It would be interesting to search for the HB in the outer portion ($R_p>40\\arcsec$) of the ACS field, where young stars should be quite rare and the old component should be dominant.  \\item A possible correlation was noted between the asymmetric distributions of the young stars and of the neutral Hydrogen. Since the young population is concentrated in the innermost part of the galaxy, it can be traced in much deeper detail with the ACS data. This, in turn, prompts a more stringent comparison with the \\HI distribution.  \\item We identified a candidate cluster located near the center of the galaxy and within the sheet of young stars (Cluster~1, C1 hereafter). The integrated spectrum of the source, shown in Pap-I, shows strong H$_{\\beta}$ in absorption, suggesting a young age. With the high spatial resolution ACS image, we can hope to ascertain the real nature of C1 and look for other conspicuous clusters, if any.  \\end{itemize}  There is no reason to repeat the excellent analysis of J11, hence we focus strictly on the scientific cases listed above.  In the following, we adopt the parameters of VV124 derived in Pap-I (see their Tab.~1), if not stated otherwise. In particular, we adopt $D=1.3$~Mpc (in good agreement with J11) and $E(B-V)=0.015$ \\citep[from][]{ebv}. At this distance $1\\arcsec$ corresponds to 6.3~pc, $1\\arcmin$ to 378~pc. We also adopt the de-projected coordinate system X,Y of Pap-I, centered on the center of VV124 and rotated such to let the X axis coincide with the major axis of the galaxy. Finally, we use the elliptical radial distance from the center of the galaxy $r_{\\epsilon}$. To give an example, stars fulfilling the condition $r_{\\epsilon}\\le 1\\arcmin$ are enclosed within an ellipse whose major axis coincides with the major axis of VV124 and has $a=1\\arcmin$ and $\\epsilon=0.44$ (see Fig.~\\ref{map}, below). The normal (projected) radial distance from the center of the galaxy will be denoted with $R_p$. All theoretical models used for comparison with observed CMDs are from the (canonical, solar-scaled) BASTI set \\citep{basti,perci},  in the native ACS- Wide Field Channel (WFC) VEGAMAG system\\footnote{\\tt http://193.204.1.62/index.html}.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{map_lr.ps} \\caption{{\\em Upper panel:} stars detected in the ACS images (dark dots) are superposed to the stars from the LBT sample of Pap-I (grey dots) in a projected coordinate system centered on the center of VV124 and rotated to have the X axis coinciding with the major axis of the galaxy (see Pap-I). Only stars brighter than $I=26.5$ from the ACS sample have been plotted, for clarity. {\\em Lower panel:} zoomed view of the central part of the ACS sample in the same coordinate system. Two concentric ellipses of semi-major axis $a=1.0\\arcmin$ and $a=1.5\\arcmin$, having the same orientation and ellipticity of the galaxy isophotes (see Pap-I) are overplotted (in white and in black, respectively). The white circle has a radius $R_p=40\\arcsec$, enclosing the region studied by J11. Blue stars lying on the young Main Sequence ($-0.4\\le V-I\\le 0.0$ and $I<25.5$, see Fig.~\\ref{comcm}) are plotted as red filled circles, the brightest ones (having $I\\le 23.0$) are encircled in orange.} \\label{map} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{CFplot_lr.ps} \\caption{Results of the artificial stars experiments. {\\em Upper panels: } Difference between the {\\em output} and {\\em input} magnitudes of the artificial stars as a function of F814W magnitude. {\\em Lower panels:} Completeness as a function of F606W magnitude in the inner ($r\\le 1.0\\arcmin$; left panel) and outer ($r> 1.0\\arcmin$; right panel) regions of the galaxy. Red continuous curves shows the completeness for stars in the color range of the RGB and the RC, blue dashed curves refer to stars in the color range of the old HB and the YMS (see Fig.~\\ref{comcm}, below). The blue curves start at $F606W=25.0$ because in that color range we produced only artificial stars fainter than this limit, focussing to the study of the completeness in the blue HB region.} \\label{CFplot} \\end{figure}    The plan of the paper is as follows: in Sect.~\\ref{obs} we describe the observational material and the data reduction process; in Sect.~\\ref{bhb} we present the clear detection of the extended HB in the external regions of the ACS field and we discuss the radial distribution of the various stellar populations of VV124. Sect.~\\ref{sec:HI} is devoted to a brief discussion of the correlation between the spatial distribution of the \\HI and of the young stars, and in Sect.~\\ref{clus} we study the candidate luminous star clusters identified here and in Pap-I. Finally we briefly summarize and discuss our main results in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  We have presented a new reduction and analysis of archival HST-ACS images of the isolated dwarf galaxy VV124. A detailed study of the SFH in the innermost regions of the galaxy, from the same data,  has been already presented by another group (J11). Here we have focused our attention on a few specific issues that have been raised in Pap-I and were not addressed in the J11 paper.  \\begin{itemize}  \\item We provide a clear detection of an extended Horizontal Branch population. It has been shown that HB stars have a spatial distribution which is less concentrated than the RC population. This strongly supports the idea that very old and metal poor stars become more and more dominant at larger distances from the center of VV124.  \\item The near coincidence of the asymmetric sheet of young stars near the center of VV124 and the density peak of the (also asymmetric)  associated \\HI cloud strongly support the idea that the complex kinematics and velocity gradients observed in the gas in Pap-I are most likely associated with the dynamics of the star formation episode. It is interesting to note that some of the isolated dwarfs in the simulations by \\citet{S11} lose all their \\HI at very early times and are able to re-accrete some of the expelled gas at later times. This scenario would be fully consistent with the observed properties of the \\HI in VV124 and would provide a natural explanation for the long period of quiescence between the early phase of star formation and the very recent sprout of activity that originated the sparse population of young MS stars currently observed (see J11).  \\item We confirm that the candidate cluster identified in Pap-I (C1) is a genuine young massive star cluster \\citep[see][and references therein]{all}. It has an age of $\\sim 250$~Myr and a total mass of $\\simeq 1.2\\times10^4~M_{\\sun}$. The cases of such low-mass early-type dwarf galaxies with massive clusters are quite rare in the Local Group (M98). We identified another likely cluster (C2), significantly younger (age$\\le 30$~Myr) and less massive ($M\\le 3.3\\times10^3~M_{\\sun}$) than C1.  \\end{itemize}"
    },
    "1107/1107.5695_arXiv.txt": {
        "abstract": " ",
        "introduction": "It has seemingly been taken for granted that in 1929 Hubble discovered exactly what Friedmann predicted several years before, in 1922 (see, for example, \"The brief history of time\" by Stephen Hawking [1], a book issued in a number of copies which is much larger than the total number of copies of all other books on cosmology ever published). A book by Alexander Sharov and Igor Novikov declares in its title: \"Edwin Hubble, the Discoverer of the Big Bang Universe\" [2]. Einstein was among the first physicists and astronomers who adopted or shared this view (but not Friedmann who died in 1925).  A non-traditional point was however made by Steven Weinberg in \"The First Three Minutes\"[3]: \"Actually, a look at Hubble's data leaves me perplexed how he could reach such a conclusion -- galactic velocities seem almost uncorrelated. In fact, we would not expect any neat relation of proportionality between velocity and distance for these 18 galaxies -- they are all much too close, none been further than the Virgo Cluster. It is difficult to avoid the conclusion that... Hubble knew the answer he wanted to get.\"  In [1-3], Lema$\\hat{\\imath}$tre is not given any credit for the linear velocity-distance relation (the very his name cannot be found in [1]), though the discovery of the relation is explicitly reported in his 1927 paper long known to many cosmologists. The observational data he used (especially as presented in a velocity-distance diagram [4]) look somewhat more scattered than that in Hubble's diagram of 1929. So Weinberg's challenge extends equally to Lema$\\hat{\\imath}$tre's result as well.  Contrary to Weinberg, Alan Sandage, who was Hubble's successor in observational cosmology on Mount Wilson and Mount Palomar, accepted Hubble's law for local galaxies as an important well-established empirical fact. He had however doubts about its cosmological interpretation [5-7].  The recent debate on the history of Hubble's law [4,8-11] revives interest to earlier controversial views [1-3,5-7] and rises again a question on the essence of the matter: What was actually discovered -- if any -- in 1927-29? ",
        "conclusions": "In 1927, Lema$\\hat{\\imath}$tre discovered dark energy and Hubble confirmed this in 1929.   \\vspace{0.3cm}  I thank G. Byrd, Yu. Efremov, I. Karachentsev, D. Makarov, P. Teerikorpi, M. Valtonen for collaboration and many useful discussions."
    },
    "1107/1107.1280_arXiv.txt": {
        "abstract": "The quasar SDSS~J133401.39+331534.3 at $z=2.426$ is found to be a two-image gravitationally lensed quasar with the image separation of $0\\farcs833$. The object is first identified as a lensed quasar candidate in the Sloan Digital Sky Survey Quasar Lens Search, and then confirmed as a lensed system from follow-up observations at the Subaru and University of Hawaii 2.2-meter telescopes. We estimate the redshift of the lensing galaxy to be $0.557$ based on absorption lines in the quasar spectra as well as the color of the galaxy. In particular, we observe the system with the Subaru Telescope AO188 adaptive optics with laser guide star, in order to derive accurate astrometry, which well demonstrates the usefulness of the laser guide star adaptive optics imaging for studying strong lens systems. Our mass modeling with improved astrometry implies that a  nearby bright galaxy $\\sim 4''$ apart from the lensing galaxy is likely to affect the lens potential. ",
        "introduction": "\\label{sec:intro}  Since the discovery of the first gravitationally lensed quasar \\citep{walsh79}, it has convincingly been demonstrated that lensed quasars provide insights into various research fields in astrophysics and cosmology, as well as being a unique tool for studying the dark universe. For instance, one can directly measure the structure and substructure of lensing objects, including the distribution of dark matter, through detailed observations and analysis of lensed systems \\citep[e.g.,][]{treu04,rusin05,chiba05}. Magnification by gravitational lensing greatly enhances our ability to study quasar host galaxies at high-redshifts \\citep[e.g.,][]{peng06}. Moreover, the statistics of lensed quasars, as well as measurements of time delays between lensed images, provide independent constraints on cosmological parameters \\citep[e.g.,][]{chae02,oguri07,oguri08,suyu10}.  High resolution imaging such as that provided by Adaptive Optics (AO) is crucial for accurately characterizing lensed quasar systems. This is not only the case for the small-separation systems ($\\sim1''$), where accurate astrometry (prone to the effect of atmospheric seeing) is necessary for constraining mass models of the lensing galaxies \\citep[e.g.,][]{sluse08}, but also for the larger separation systems, where adaptive optics imaging can identify lens substructures and faint extended lensed features such as arcs \\citep[e.g.,][]{mckean07}.  The Sloan Digital Sky Survey Quasar Lens Search \\citep[SQLS;][]{oguri06,oguri08,inada08,inada10} has been one of the most successful strong lens surveys conducted to date. SQLS relies on the large homogeneous sample of spectroscopically-confirmed quasars from the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}. By identifying lensed quasar candidates using the color and morphological information from the SDSS imaging data, SQLS has discovered more than 40 new lensed quasars so far. Together with previously known lenses which were re-observed by the SDSS, the SQLS sample consists of $\\sim 60$ lensed quasars, constituting roughly half of all lensed quasars discovered to date.  However, a disadvantage of SQLS, like other ground-based optical strong lens surveys, is its poor capability of identifying subarcsecond (image separation $\\theta<1''$) lensed systems. This is because typical seeing sizes of ground-based optical imaging observations are $\\sim 1''$, which makes it difficult to resolve multiple components of subarcsecond systems. Indeed, discoveries and confirmations of subarcsecond lensed quasars have been made mostly by radio surveys \\citep[e.g.,][]{browne03} or by using the Hubble Space Telescope \\citep[e.g.,][]{reimers02}, with few exceptions from ground-based surveys \\citep[e.g.,][]{castander06,blackburne08}.  In this paper, we report the discovery of the first new subarcsecond (two-image) lensed quasar discovered in the course of SQLS, SDSS~J133401.39+331534.3 (hereafter SDSS~J1334+3315). After briefly describing the lens candidate selection from the SDSS data (\\S\\ref{sec:sdss}), we present our imaging and spectroscopic follow-up observations (\\S\\ref{sec:obs}). We further observe this system with the Subaru Telescope laser guide star adaptive optics (LGS$+$AO188), for robust detection as well as characterization of the lensing galaxy (\\S\\ref{sec:lgsao}, \\S\\ref{sec:photoz}). The accurate astrometry obtained with the LGS$+$AO188 is then used for gravitational lens mass modeling (\\S\\ref{sec:mass}). We summarize our results in \\S\\ref{sec:summary}. We assume the concordance cosmology with $H_0=70 \\ \\mathrm{km}^{-1}\\ \\mathrm{s}^{-1}\\ \\mathrm{Mpc}^{-1}$, $\\Omega_m=0.27$ and $\\Omega_\\Lambda=0.73$ throughout this paper. ",
        "conclusions": "\\label{sec:summary}  We have reported the discovery of the subarcsecond ($\\theta=0\\farcs833$) gravitationally lensed quasar SDSS~J1334+3315. The system has been identified as a lensed candidate from the SDSS data, using the standard SQLS candidate selection algorithm, despite the image separation being smaller than the criterion for which the morphological selection is complete. Our follow-up observations at the Subaru and UH88 telescopes have confirmed the system to be a real strong lensing event. It consists of two images of a quasar at $z=2.426$, lensed by a foreground galaxy. From the colors as well as absorption lines in the quasar spectra, we infer the redshift of the lensing galaxy to be $z=0.557$. In particular, we have obtained high-resolution near-infrared images of this lensed quasar system using the Subaru LGS$+$AO188 system, with a better resolution in the $K'-$band than would be provided by the Hubble space telescope in the equivalent band. The images clearly reveal the presence of the lensing galaxy in between the quasar images, and enable us to derive very accurate relative astrometry, as well as shapes of galaxies. In particular, the higher resolution made possible to identify errors in the non-adaptive optics astrometry, caused by the lens galaxy and one of the quasar images being blended, and thus made possible a more accurate mass modeling. We have found that the image configuration can well be reproduced with the standard mass models. We have pointed out that the nearby galaxy G1, which is located $<5''$ from the lensing galaxy and is likely associated physically with the lensing galaxy, may affect the lens potential. This lensed system has the smallest image separation among the newly discovered SQLS lenses.  This work represents the first observation and analysis of any gravitationally lensed quasar system with the Subaru LGS$+$AO188. Our results demonstrate the usefulness of laser guide star adaptive optics imaging observations, in particular that of the Subaru Telescope LGS$+$AO188 imaging, which has a much improved sky coverage, for the study of strong lens systems. As explicitly shown, the high-resolution imaging enabled by the adaptive optics system is crucial for accurate and robust measurements of the position and shape of the lensing galaxy, particularly for small-separation lensed quasars as presented in the paper. Indeed we are conducting an imaging survey of SQLS lenses with Subaru LGS$+$AO188, aiming to derive accurate astrometry and to detect lensed quasar host galaxies, which will further increase the value of the SQLS lensed quasar sample."
    },
    "1107/1107.4977_arXiv.txt": {
        "abstract": "We investigate the {\\it Fermi} LAT $\\gamma$-ray and 15 GHz VLBA radio properties of a joint $\\gamma$-ray- and radio-selected sample of AGNs obtained during the first 11 months of the {\\it Fermi} mission (2008 Aug 4 - 2009 Jul 5). Our sample contains the brightest 173 AGNs in these bands above declination $-30^\\circ$ during this period, and thus probes the full range of $\\gamma$-ray loudness ($\\gamma$-ray to radio band luminosity ratio) in the bright blazar population.  The latter quantity spans at least four orders of magnitude, reflecting a wide range of spectral energy distribution (SED) parameters in the bright blazar population.  The BL Lac objects, however, display a linear correlation of increasing $\\gamma$-ray loudness with synchrotron SED peak frequency, suggesting a universal SED shape for objects of this class. The synchrotron self-Compton model is favored for the $\\gamma$-ray emission in these BL Lacs over external seed photon models, since the latter predict a dependence of Compton dominance on Doppler factor that would destroy any observed synchrotron SED peak - $\\gamma$-ray loudness correlation. The high-synchrotron peaked (HSP) BL Lac objects are distinguished by lower than average radio core brightness temperatures, and none display large radio modulation indices or high linear core polarization levels. No equivalent trends are seen for the flat-spectrum radio quasars (FSRQ) in our sample. Given the association of such properties with relativistic beaming, we suggest that the HSP BL Lacs have generally lower Doppler factors than the lower-synchrotron peaked BL Lacs or FSRQs in our sample. ",
        "introduction": "\\label{s:intro}  The successful launch of the \\fermi {\\it Gamma-Ray Space Telescope} in 2008 has brought about a new era in our understanding of blazars, which dominate the extragalactic sky at high energies. Because of their highly variable fluxes and spectral energy distributions (SEDs), blazar samples are typically subject to large biases, making it difficult to study their demographics. With the nearly continuous all-sky monitoring capabilities of {\\it Fermi's} Large Area Telescope (LAT), however, it is now possible to construct well-defined samples that can be used to investigate the wide range of jet properties in these powerful AGNs (e.g., \\citealt{1LAC,2009ApJ...707L..56K}).  One of these properties that has been of considerable interest since the era of the {\\it Compton Gamma Ray Observatory} (CGRO) in the 1990s is \\gr loudness, or in other words, why only a particular small subset of known AGNs ($\\sim 100$; \\citealt{Hartman99}) were detected by the CGRO's EGRET telescope. Considerable evidence has been presented by many researchers (e.g., \\citealt{1995MNRAS.273..583D,2cmPaperIII, 2cmPaperIV, J01, Taylor07}) supporting the idea that relativistic Doppler boosting has a large impact on AGN \\gr emission, but lingering questions regarding the roles of the flaring duty cycle and the AGN spectral energy distribution remain. The superior sensitivity and full-time survey operation mode of \\fermi has now provided substantial insight into these issues. With the release of the 1FGL catalog \\citep{1FGL}, the strong impact of SED characteristics on the fainter \\gr AGN population was realized, as the sky at these levels becomes dominated by high-synchrotron peaked BL Lac objects. At the same time, the predominant association of \\fermi LAT sources with flat-spectrum radio quasars (FSRQ) and BL Lacs (blazars) has established Doppler boosting as the primary factor in determining \\gr loudness in the brightest AGNs.  In this paper, we follow up on previous analyses of bright blazars that were based on the initial 3 month LAT data set presented by \\cite{LBAS}. These studies established several important AGN radio/$\\gamma$-ray connections using quasi-simultaneous VLBA observations, namely that the \\gr photon flux correlates with the parsec scale radio flux density \\citep{MF2, Arshakian2011}, and that the jets of the LAT-detected blazars have higher-than-average apparent speeds \\citep{MF1}, larger apparent opening angles \\citep{Pushkarev2010}, more compact radio cores \\citep{MF2}, strong polarization near the base of the jet \\citep{2011ApJ...726...16L}, and higher variability Doppler factors \\citep{MF4}. In addition, AGN jets have been found to be in a more active radio state within several months of the LAT-detection of their strong \\gr emission \\citep{MF2}, which was subsequently confirmed by \\cite{2010ApJ...722L...7P}.  With the release of the 1st LAT AGN catalog (1LAC; \\citealt{1LAC}) based on the initial 11 months of \\fermi data, it is now possible to investigate the impact of Doppler beaming and SED characteristics on AGN \\gr loudness using larger, more complete samples and better statistics. Here we present a joint analysis of \\fermi and VLBA 15 GHz radio properties of the brightest radio and \\gr AGN in the northern sky, based on data from the LAT instrument, flux density measurements from the OVRO and UMRAO radio observatories, and the MOJAVE VLBA program \\citep{MOJAVE_V}. In particular we examine the differences in the SED and \\gr properties of BL Lac objects with respect to FSRQs, and the relative role of relativistic beaming on their \\gr loudness. Several complementary studies will examine the connection between \\gr emission and superluminal speeds (Kadler et al., in prep.), detailed SED parameters (Chang et al., in prep.), and radio jet activity level (Lister et al., in prep.).  Throughout this paper, we use a $\\Lambda$CDM cosmological model with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$ \\citep{Komatsu09}. ",
        "conclusions": "\\subsection{Redshift Distributions}  \\begin{figure}[t] \\centering \\includegraphics[width=0.47\\textwidth,trim=0cm 0cm 0cm 0cm,clip]{zhistograms.ps} \\caption{\\label{zhistograms} Left panel: redshift distribution of the \\grd selected 1FM sample. The full sample is represented by the solid line, and the BL Lac objects are shaded. There is one radio galaxy (J0319+4130 = 3C 84) at $z = 0.0176$. Right panel: redshift distribution for the radio-selected 1FM matching sample. There are four radio galaxies in the sample, all in the first ($z < 0.25$) bin.} \\end{figure}  The redshift data on our AGNs are incomplete (see Appendix), with missing values for 4 sources in the radio-selected sample, and 22 sources in the \\grd selected sample (the sources J0050$-$0929 and J0818+4222 are common to both samples).  In Figure~\\ref{zhistograms} we plot the redshift distributions for our samples. The redshifts range from $z = 0.00436$ to $z = 3.396$, and the distributions are generally peaked between $z=0.5$ and $z=1$. Kolmogorov-Smirnov (K-S) tests do not reject the null hypothesis that the \\grd selected and radio-selected samples are drawn from the same parent redshift distribution, even when the sources in common to both samples are excluded (D = 0.20, probability = 0.27). We find no statistical differences in the redshift distributions of the non-LAT detected versus LAT-detected AGNs in the combined samples (D = 0.16, probability = 0.49).  With respect to the redshifts of the quasars in the two samples, the K-S test suggests a marginal statistical difference in their distributions (D = 0.079, probability = 0.96). There are an insufficient number of radio galaxies to perform any statistical tests on them (there are 4 in the radio-selected sample; one of these is also in the \\grd selected sample). The overall redshift distribution of the \\grd selected sample has an additional peak at low redshift, due to the presence of at least 9 high synchrotron peaked (HSP) BL Lacs that are not in the radio selected sample (8 additional HSP BL Lacs lack redshift information). These objects also bring the overall fraction of BL Lacs up to 35\\% in the \\grd selected sample, as compared to only 13\\% for the radio-selected sample.   Because of the similarities in the properties and redshift distributions of the \\gr and radio-selected samples, for the remainder of this paper we will no longer distinguish between them, referring instead to the joint sample of 173 AGNs.  \\subsection{\\gr Loudness and Synchrotron Peak Frequency}  \\begin{figure}[t] \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 2cm,clip]{lgam_vs_lradio.ps} \\caption{\\label{lgam_vs_lradio} Plot of average \\gr luminosity versus median VLBA 15 GHz radio luminosity.  The filled circles represent BL Lac objects, with the high synchrotron peaked ones in orange and others in blue. The open circles represent quasars and the green diamonds radio galaxies. The arrows represent upper limits based on the 11-month LAT data.  The dashed line represents the 1:1 luminosity ratio line. } \\end{figure}  A primary goal of our study is to examine the range of \\gr loudness ($G_r$) present in the bright blazar population, and its dependence on other AGN jet properties.  Since we have obtained data in several complete regions of the \\grd radio plane (i.e., \\grd bright/radio-faint; \\grd faint/radio-bright; \\grd bright/radio-bright) we can be assured of sampling the largest possible range of $G_r$ in the brightest northern-sky blazars. Future studies of the \\grd weak/radio-weak region will be important for verifying whether the trends we identify here extend to the fainter blazar population.  In Figure~\\ref{lgam_vs_lradio} we plot \\gr luminosity against 15 GHz VLBA luminosity.  Despite our use of an average \\gr luminosity over an 11-month period, the linear relationship for the non-censored data has only moderate scatter (0.6 dex). A linear regression fit to the non-censored data yields $\\log{L_{\\gamma}} = (0.92 \\pm 0.05) \\log{L_R} + 6.4 \\pm 2$.  The $G_r$ values, which reflect the perpendicular distance of the data points from the dashed 1:1 line, span nearly 4 orders of magnitude, from below 3 to $\\sim 15000$. A clear division between the HSP and lower-synchrotron-peaked BL Lacs is evident, with the former having higher \\gr loudness ratios.  None of the radio galaxies are significantly \\grd loud, with ratios all below 65.  The quasars and BL Lacs have significantly different $G_r$ distributions (Figure~\\ref{rg_histograms}), with the former peaking at $G_r \\simeq 10^3$ and the latter peaking above $10^{3.5}$. There is a substantial population of quasars with $G_r$ values below 100, while all of the BL Lacs (with the exception of J0006$-$0623) have \\Fermi associations  and $G_r > 60$.  The Peto and Peto modification of Gehan's Wilcoxon two-sample test for censored data rejects the null hypothesis that the quasar and BL Lac $G_r$ values come from the same parent population at the 99.99\\% confidence level.  \\begin{figure}[t] \\centering \\includegraphics[width=0.45\\textwidth,trim=0cm 0cm 0cm 0cm,clip]{rg_histograms.ps} \\caption{\\label{rg_histograms} Distribution of \\gr to radio luminosity ratio for BL Lacs (top panel) and quasars (bottom panel). Upper limit values for AGNs with no \\fermi 1LAC catalog associations are indicated by the dashed lines.  } \\end{figure}   These differences are reflected in Figure~\\ref{rg_vs_sedpeak}, which shows \\gr loudness plotted against the synchrotron SED peak frequency. The BL Lacs show a roughly linear correlation of the form $\\log{G_r} = (0.40\\pm 0.06)\\log{\\nu_s} - 2.9 \\pm 0.9$, with a scatter of 0.5 dex, while the quasars show no trend.  It is apparent that the BL Lacs have a higher mean \\gr loudness value because many of them have synchrotron peaks above $\\sim 10^{15}$ Hz.  Since the fixed radio bandpass is always located below the synchrotron peak, if we compare two BL Lacs with identical SED shapes but different synchrotron peak locations, the high synchrotron peaked BL Lac will have a lower radio flux density, and thus a higher \\gr loudness value. Figure~\\ref{sed_peak_vs_fluxradio} shows this broad trend for the BL Lacs, with the HSP jets having generally lower radio flux densities than the LSPs.  A similar spectral index effect also occurs as the high energy SED peak moves in tandem through the \\fermi LAT band as the synchrotron peak frequency increases. This is manifested in the strong correlation seen between the \\gr photon spectral index $\\alpha_G$ and synchrotron peak frequency for the 1FGL blazars, as described by \\cite{1LAC}. In Figure~\\ref{rg_vs_energyindex} we plot \\gr loudness against photon spectral index. Again we see a good (even tighter) linear correlation for the BL Lac objects, and no trend for the quasars. A regression fit to the BL Lacs, omitting the outlier source J0825+0309, gives $\\log{G_r} = (-2.3 \\pm 0.2)\\alpha_G + 7.8 \\pm 0.5$, with a scatter of 0.3 dex.  \\begin{figure}[t] \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.7cm,clip]{rg_vs_sedpeak.ps} \\caption{\\label{rg_vs_sedpeak}  \\gr to radio luminosity ratio $G_r$ versus synchrotron SED peak frequency. The red filled circles represent BL Lac objects, the open circles quasars, the green diamonds radio galaxies, and the purple crosses optically unidentified objects. The arrows denote upper-limits. The BL Lac objects show a linear trend of increasing \\gr loudness with SED peak frequency, while no trend exists for the quasars. } \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.7cm,clip]{sed_peak_vs_fluxradio.ps} \\caption{\\label{sed_peak_vs_fluxradio} 15 GHz VLBA flux density versus synchrotron SED peak frequency.  The red filled circles represent BL Lac objects, the open circles quasars, the green diamonds radio galaxies, and the purple crosses optically unidentified objects. } \\end{figure}  The continuous trend from LSP to HSP BL Lacs in Figures~\\ref{rg_vs_sedpeak} and \\ref{rg_vs_energyindex} is noteworthy, since it implies a relatively narrow intrinsic range of variation in the SED shapes of the brightest BL Lac objects. Broadly speaking, there are three aspects of  a SED that can affect its measured \\gr loudness parameter. These are the relative positions of the synchrotron and high energy peaks with respect to the fixed \\gr and radio bands, their relative luminosities (often referred to as the Compton dominance), and the width and shape of each peak. If we take the simplest case of both peaks having equal luminosity and identical parabolic forms in $\\nu F_\\nu$ --$\\nu$ space, then we would expect to have \\begin{equation} \\log{G_r} = C_1 \\left(\\log{ \\nu_h \\over \\nu_\\gamma}\\right)^2 - C_2\\left(\\log{ \\nu_s \\over \\nu_r}\\right)^2, \\end{equation} where $\\nu_\\gamma$ and $\\nu_r$ are the frequencies of the LAT \\gr and VLBA radio bands, $\\nu_h$ and $\\nu_s$ are the frequencies of the high energy and synchrotron peaks, and $C_1$ and $C_2$ are parameters that determine their respective widths.  If both SED peaks have identical parabolic shape ($C_1$ = $C_2$) and the entire SED is then shifted to a higher frequency, such that the peak separation $\\log{\\nu_h}-\\log{\\nu_s} = C_3 $ remains constant, then we would expect to find a linear relation of the form $\\log{G_r} = a \\log{\\nu_s}$, with slope \\begin{equation} a = 2 C_1 \\left(\\log{\\nu_r \\over \\nu_\\gamma} + C_3\\right). \\end{equation}  \\begin{figure}[t] \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.1cm,clip]{rg_vs_energyindex.ps} \\caption{\\label{rg_vs_energyindex} Plot of \\gr to radio luminosity ratio $G_r$ versus \\gr photon spectral index. The filled circles represent BL Lac objects, with the high synchrotron peaked ones in orange and others in blue. The open circles represent quasars, and the green diamonds radio galaxies. The BL Lac objects show a log-linear trend of decreasing \\gr loudness with photon spectral index, while no trend exists for quasars. } \\end{figure}  From compilations of observed blazar SEDs (e.g., \\citealt{abdo_sed,Chang_thesis}) we know that in actuality, SED peak shapes are only approximately parabolic, and that there exists a range of $C$ parameter values among the population.  These factors would tend to distort any trend between $\\log{\\nu_s}$ and $\\log{G_r}$ from the simple linear one described here.  Furthermore, an intrinsic range of Compton dominance parameter would likely destroy any linear relation completely.  The fact that we see a scatter of only 0.5 dex for the BL Lac objects therefore implies that the SEDs of the brightest AGNs of this class must have relatively similar shapes, at least much more so than the quasars, which show no $\\nu_s$-$G_r$ correlation. Our results are corroborated by a recent study of the 1LAC by \\cite{2011arXiv1106.5172G}, who defined a ''Compton efficiency'' parameter as the ratio of the high-energy (inverse Compton) SED peak luminosity to 8 GHz radio VLA core luminosity. They found a similar trend of higher Compton efficiency with increasing synchrotron peak frequency for BL Lac objects, but no trend for FSRQs.  So far in this discussion we have omitted the possible effects of relativistic beaming on the SED. For the simple case of the same Doppler factor in both the radio and \\grd emitting regions, the entire SED should be blue-shifted by the Doppler factor, and the apparent luminosity of both peaks will be increased by Doppler boosting. Models which attribute the high energy peak to inverse Compton scattering of external seed photons by relativistic electrons in the jet predict a higher Doppler boost in $\\gamma$-rays, because of the additional Lorentz transformation between the seed photon and jet rest frames \\citep{Dermer1995}.  In this case, when considering a jet at smaller viewing angle, the resulting increase in Doppler factor boosts the luminosity of the high energy peak to a level much higher than the synchrotron peak, thereby increasing the observed Compton dominance and \\gr loudness. If the seed photons are internal to the jet, for a single-zone synchrotron self-Compton (SSC) model relatively equal boosting is expected in both regimes; thus $G_r$ in this case is much less sensitive to Doppler boosting. The fairly good linear $G_r$-$\\nu_s$ correlation for the BL Lacs therefore favors the SSC process as the dominant emission mechanism in this class of blazars. This is in general agreement with the conclusions of recent studies which have modeled the SEDs of {\\it Fermi}-detected blazars with detailed synchrotron and inverse-Compton emission models \\citep{abdo_sed}. It should be possible to investigate this issue in much greater depth when more detailed information on the SED parameters of our full sample can be obtained.   \\subsection{Parsec-Scale Radio Jet Properties} \\subsubsection{Core Brightness Temperature}  Nearly all of the AGNs in our sample have a parsec-scale radio jet morphology that is dominated by a bright, flat-spectrum core, which is often unresolved or barely resolved in our mas-scale VLBA images. At our observing frequency of 15 GHz, this core typically represents the region where the jet becomes optically thick, with the true jet nozzle being located upstream \\citep{2011arXiv1103.6032S}. The brightness temperature of the core component in our VLBA images measures the compactness of the radio jet emission, and has been previously shown to be correlated with indicators of relativistic beaming, such as superluminal apparent speed \\citep{Homan_etal06} and radio flux density variability \\citep{Tingay01, Hovatta09}.  In Figure~\\ref{tb_vs_sedpeak} we plot core brightness temperature against synchrotron SED peak frequency. The main visible trend is that the HSP BL Lac radio cores tend to be less compact than those of the other AGNs in the sample. We discuss the possible ramifications of this trend on beaming and jet velocity stratification models for HSP BL Lacs in \\S~\\ref{HSP}.   \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.7cm,clip]{tb_vs_sedpeak.ps} \\caption{\\label{tb_vs_sedpeak} Radio core brightness temperature at 15 GHz versus synchrotron SED peak frequency. The filled circles represent BL Lac objects, with the high synchrotron peaked ones in orange and others in blue.  The open circles represent quasars, the green diamonds radio galaxies, and the purple crosses optically unidentified objects. The arrows denote lower limits. The radio cores of the high-synchrotron SED peak BL Lac objects tend to be less compact than the other AGNs in our sample. J1104+3812 (Mrk 421) is the only high-synchrotron peaked blazar in the sample with a high core brightness temperature. Not plotted is the unusually low-brightness temperature quasar J0957+5522 (4C +55.17) at $\\nu_s = 10^{13.77}$ Hz, $T_b = 10^{8.46}$ K.  } \\end{figure}   \\subsubsection{\\label{opangles}Apparent Jet Opening Angles}  In a previous study of the MOJAVE sample using the initial 3 months of \\fermi data, \\cite{Pushkarev2010} found a tendency for the \\grd detected blazars to have wider apparent opening angles than the non-detected ones.  Since the calculated intrinsic opening angles of the two groups were similar, they concluded that the \\grd detected jets were viewed more closely to the line of sight.  We have analyzed the apparent jet opening angles of our sample, and find that they range from 5 to 68 degrees, with a mean of 24 degrees. There is an extended tail to the distribution, with 19 jets having opening angles greater than 40 degrees.  With the exception of the quasar J0654+4514, none of the high opening angle jets are highly variable (radio modulation indices all less than 0.26). We find no statistically discernible differences in the opening angle distributions of the different optical or SED classes.  We do find a correlation between \\gr loudness and apparent opening angle, however the relationship is non-linear (Fig.~\\ref{rg_vs_opangle}). All of the AGNs in the high-opening angle tail ($>40$ deg.) of the distribution are significantly \\grd loud ($G_r > 100$). The apparent opening angle of a jet is related to the viewing angle and intrinsic opening angle, with smaller intrinsic angles expected for high Lorentz factor jets based on hydrodynamical considerations \\citep{2005AJ....130.1418J}. The high opening angle jets in our sample are a mixture of BL Lac objects and FSRQ, with a range of synchrotron SED peak frequencies. With jet kinematic information on the sample from the MOJAVE program it will be possible to investigate whether these particular jets are viewed unusually close to the line of sight, or have atypically large intrinsic opening angles.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.1cm,clip]{rg_vs_opangle.ps} \\caption{\\label{rg_vs_opangle} \\gr to radio luminosity ratio versus apparent jet opening angle.  The filled circles represent BL Lac objects, with the high synchrotron peaked ones in orange and others in blue. The open circles represent quasars, the green diamonds radio galaxies, and the purple crosses optically unidentified objects. The AGNs with wide apparent opening angles tend to have high \\gr loudness values. } \\end{figure}  \\subsubsection{Radio Core Polarization Vectors}  We compared the direction of the linear polarization vector at the radio core position to the mean jet position angle for our sources, as described in \\S~\\ref{vlba_data}.  In some sources such as PKS 1502$+$106 \\citep{Abdo_1502+106} and PMN J0948$+$0022 \\citep{Foschini2011}, changes in the core polarization angle have been seen to occur in conjunction with \\gr flaring events, suggesting a close connection between the radio and \\gr emission regions. We do not find any correlations between the core polarization vector offset and any \\gr or SED properties for our sample. However, since the linear polarization vector angles tend to be highly variable in blazars \\citep{2005AJ....130.1418J}, a more detailed analysis would require truly simultaneous VLBA-\\fermi measurements rather than the average \\gr data that we use in our current study. Another possible reason for the lack of correlations is Faraday effects in the cores, which can rotate the observed polarization vectors. We are currently completing a VLBA rotation measure analysis of the original MOJAVE radio-selected sample \\citep{2011arXiv1108.1514H} to further investigate this effect.  \\subsubsection{Radio Core Polarization Level}  \\begin{figure}[t] \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.1cm,clip]{mcore_vs_sedpeak.ps} \\caption{\\label{mcore_vs_sedpeak} Linear fractional polarization level of VLBA radio core at 15 GHz versus synchrotron SED peak frequency. The red filled circles represent BL Lac objects, the open circles quasars, the green diamonds radio galaxies, and the purple crosses optically unidentified objects. The arrows denote upper-limits.  } \\end{figure}  In Figure \\ref{mcore_vs_sedpeak} we plot the linear fractional polarization level of the VLBA core at the reference epoch versus SED peak frequency. In general the cores of the jets are weakly polarized ($< 4$ \\%), with increasing fractional polarization levels seen downstream \\citep{MOJAVE_I}.  There are no appreciable differences in the BL Lac and FSRQ core polarization distributions, however, as we discuss in section \\ref{HSP}, the high synchrotron peak BL Lacs tend to have low core polarization levels. We find no trend between 15 GHz radio core polarization and \\gr loudness, although in the VIPS 5 GHz VLBA survey \\citep{2011ApJ...726...16L} detected core polarization more frequently in LAT-detected AGNs than in the non-LAT ones. We note that the core polarizations tend to vary over time in these jets, which can complicate such analyses. Indeed, in a preliminary investigation of the original MOJAVE radio flux-limited sample, we found the LAT-detected AGNs tended to have higher median fractional core polarization levels during the first three months of the Fermi mission, as compared to their historical average level \\citep{2010IJMPD..19..943H}. A more complete polarization analysis of our full multi-epoch MOJAVE VLBA dataset will be presented in a forthcoming study.  \\subsubsection{Radio Variability} The hallmark flux variability seen in blazars is believed to be closely related to Doppler beaming \\citep{AAH92, LV99,Hovatta09} since it can significantly heighten the magnitudes of flaring events and shorten their apparent timescales (e.g., \\citealt{Lister01}).  AGN jets have also been found to be in a more active radio state within several months from LAT-detection of their strong \\gr emission \\citep{MF2, 2010ApJ...722L...7P}. The AGNs in our sample are indeed highly variable, with 51 of 144 sources having standard deviations greater than 15\\% of their mean flux density level over an 11 month period. In their full sample of over 1000 sources, \\cite{Richards11} found the FSRQs to have significantly higher variability amplitudes than the BL Lacs. We do not see this distinction in our sample, however, most likely because ours contains a smaller proportion of high-synchrotron peaked BL Lacs. The latter tend to have moderately low radio modulation indexes, as seen in Figure~\\ref{mi_vs_sedpeak}.  \\begin{figure}[t] \\centering \\includegraphics[width=0.5\\textwidth,trim=0cm 0cm 0cm 1.7cm,clip]{mi_vs_sedpeak.ps} \\caption{\\label{mi_vs_sedpeak} Radio modulation index at 15 GHz versus synchrotron SED peak frequency. The filled circles represent BL Lac objects, the open circles quasars, the plus symbols radio galaxies, and the crosses optically unidentified objects. None of the blazars with synchrotron SED peaks  above $10^{15}$ Hz show high amplitude radio variability.} \\end{figure}   \\subsection{\\label{HSP}High Synchrotron Peaked AGN Jets and the BL Lac Blazar Class}  Previous studies of the full 1LAC catalog by the LAT team \\citep{1LAC,abdo_sed} have established that HSP BL Lacs have fundamentally different \\gr properties than the \\grd loud FSRQs. In our study we have found that the HSP BL Lacs are characterized by high \\gr to radio luminosity ratios and lower than average radio core compactness. Given these differences seen in both the radio and \\gr regimes, a fundamental question remains as to whether the lower synchrotron-peaked BL Lacs also form a jet population distinct from the FSRQs. The continuity of the trend between SED peak frequency and \\gr loudness (Figure~\\ref{rg_vs_sedpeak}) would suggest that their SED shapes are similar to the HSP BL Lacs, and thus they should be unified with them. They are also more similar to the HSPs in terms of their radio luminosity, as compared to the generally more luminous FSRQs (Figure~\\ref{lgam_vs_lradio}).  If we directly compare the radio properties of the LSP BL Lacs and LSP FRSQs in our sample, we find that the LSP BL Lacs have higher mean fractional linear polarization (4.8 $\\pm$ 2.8 versus 2.5 $\\pm$ 1.7; t=3.02, p=0.992 according to a Welch Two Sample t-test), although both classes span roughly the same range of extreme values (Fig.~\\ref{mcore_vs_sedpeak}). It is possible that beam depolarization effects may lower the mean value for the FSRQs, since they are typically at higher redshift than the BL Lacs and are thus imaged with poorer spatial resolution.  However, if we compare the LSP and HSP BL Lacs, which have similar redshift ranges, we find the latter have consistently low core polarization and modulation indices, as well as lower than average radio core brightness temperatures. Since high radio variability, core polarization, and brightness temperature are generally associated with high Doppler boosted-jets (e.g., \\citealt{Lister43GHZPR, Tingay01, Hovatta09}), the trends we find in our sample support the following scenario for the brightest \\gr and radio blazars in the sky. Because of their higher intrinsic \\gr loudness ratios and low redshifts, the HSP BL Lacs do not need to be as highly beamed to enter into flux-limited \\gr and radio samples, thus they tend to have lower Doppler boosting factors than other blazar classes. The LSP BL Lacs are less intrinsically luminous than the FSRQs, but their moderately high intrinsic \\gr loudness ratios and Doppler boosting factors combine to give them apparent \\gr and radio luminosities comparable to the fainter end of the FSRQ distribution.   The above scenario is supported by the previous results of \\cite{Nieppola2008}, who found a general trend of decreasing Doppler factor with increasing synchrotron SED peak frequency. Their sample only included blazars of the LSP and ISP class, however.  A potential test can be made with parsec-scale superluminal motion measurements, which set an upper limit on the viewing angle and a lower limit on the bulk jet Lorentz factor (see e.g., \\citealt{UP95}). One of the main unresolved problems for HSP BL Lacs has been the relatively slow apparent jet speeds detected for these objects \\citep{Piner2004, Piner2010}, despite the need for large Doppler factors to account for rapid variability seen in $\\gamma$-rays and to accurately model their SEDs (see, e.g., \\citealt{HenriSauge2006}). Several models have been put forward to address this ``Doppler factor crisis'', including decelerating flows \\citep{2003ApJ...594L..27G}, and stratified spine-sheath models \\citep{Celotti2001,2008MNRAS.385L..98T}, in which the $\\gamma$-rays originate in a high-velocity jet spine, while the radio emission (and moving blobs) are associated with a lower-Lorentz factor sheath. In this manner the radio and \\gr emission can have independent Doppler factors. As we discussed in \\S~\\ref{gamrayloudness}, however, uncorrelated beaming factors for the synchrotron and high-energy peaks would likely destroy any linear relation between SED synchrotron peak and \\gr loudness, in contrast to what we see for the BL Lacs in our survey (Figure~\\ref{rg_vs_sedpeak}). A more recent model put forward by \\cite{Lyutikov2010} involving non-steady magnetized outflows suggests the existence of different, yet correlated Doppler factors for the two SED peak regions, which can potentially preserve the $G_r$ versus synchrotron SED peak relation.  With the MOJAVE program we are currently obtaining multi-epoch VLBA measurements of all the \\grd selected radio jets in our sample, which will allow us to further investigate the connections between synchrotron SED peak frequency, apparent jet speed, jet opening angle, and Doppler factor in the brightest blazars."
    },
    "1107/1107.0908.txt": {
        "abstract": "%% Text of abstract This review provides an introduction to the generation and evolution of the Sun's magnetic field, summarising both observational evidence and theoretical models. The eleven year solar cycle, which is well known from a variety of observed quantities, strongly supports the idea of a large-scale solar dynamo. Current theoretical ideas on the location and mechanism of this dynamo are presented.  The solar cycle influences the behaviour of the global coronal magnetic field and it is the eruptions of this field that can impact on the Earth's environment. These global coronal variations can be modelled to a surprising degree of accuracy. Recent high resolution observations of the Sun's magnetic field in quiet regions, away from sunspots, show that there is a continual evolution of a small-scale magnetic field, presumably produced by small-scale dynamo action in the solar interior.  Sunspots, a natural consequence of the large-scale dynamo, emerge, evolve and disperse over a period of several days. Numerical simulations can help to determine the physical processes governing the emergence of sunspots. We discuss the interaction of these emerging fields with the pre-existing coronal field, resulting in a variety of dynamic phenomena. ",
        "introduction": "\\label{sec:introduction} The Sun's magnetic field is generated through dynamo action in the solar interior. While the key physical processes involved with this dynamo are hidden from view, the consequences of the resulting magnetic field can be seen in the solar atmosphere. The large-scale magnetic field is characterised by the appearance of sunspots, regions of intense magnetic field, the number of which varies with a cycle of approximately eleven years. In addition, there is a small-scale magnetic field that varies on a much shorter timescale. The appearance of sunspots and the continual motion of the small-scale field ensure that the magnetic field of the solar atmosphere is very dynamic, with many solar phenomena owing their existence to this evolving field. The aim of this article is to describe the present theoretical ideas behind various aspects of the solar magnetic field. In particular, we shall look at current ideas for the dynamo mechanism(s) responsible for the generation of large- and small-scale fields, and will then consider the dynamical behaviour of the surface manifestations of these fields, how they might burst through the solar surface and interact with magnetic field in the solar atmosphere. We have tried to address an audience familiar with aspects of the Earth's magnetic field, but maybe less so with those of the Sun. Thus it is perhaps appropriate first to present a short description of the basic properties of the solar interior and the solar atmosphere.  \\subsection{Properties of the solar interior} The radius of the Sun is $\\Rsun = 696$\\,Mm, or equivalently $69\\,600$\\,km. From the theory of stellar structure, it follows that the interior of the Sun has three physically distinct regions. At the centre, extending to approximately $0.25 \\Rsun$, is the core, in which nuclear reactions fuse hydrogen into helium, with the associated release of energy. From $0.25 \\Rsun$ to $0.71 \\Rsun$, through what is known as the radiative zone, energy is transferred by the emission and absorption of photons. The outer $30\\%$ by radius of the Sun (which, however, constitutes only $2\\%$ of the solar mass) is the convection zone; here, the increase in opacity leads to a temperature gradient sufficiently steep to trigger convective motions. Recently, in only the past thirty years or so, it has become possible actually to determine properties of the solar interior via observation and inversion of the sound waves propagating through the interior --- a field known as \\textit{helioseismology}. Of particular interest for theories concerning the generation of the solar magnetic field is the deduction of the solar interior rotation rate, $\\Omega$. Although, through observations of surface magnetic features, it has been known for hundreds of years that the solar surface rotates differentially, with the equator rotating faster than the poles, with periods of approximately 25 and 34 days respectively, the nature of the internal rotation rate remained a matter for theoretical speculation. The actual internal rotation rate revealed by helioseismology turned out to be rather surprising \\citep{Schou_etal_98}. The differential rotation in latitude of the surface is maintained throughout the convection zone, with very little radial variation; the radiative zone is in solid body rotation, with an angular velocity comparable with that of the surface at approximately $30^\\circ$ latitude. The abrupt change in angular rotation between the convection and radiative zones is then accommodated in a very thin shear layer, known as the \\textit{tachocline}; $\\partial \\Omega / \\partial r > 0$ at low latitudes, $\\partial \\Omega / \\partial r < 0$ at higher latitudes. Although very thin in extent, occupying just a few per cent of the solar radius, the existence of the tachocline is of great significance both in determining the long-term evolution of the Sun and also for its possible role in the generation of the solar magnetic field; all currently known aspects of tachocline dynamics are covered in the recent volume edited by \\cite{HRW_07}.  Magnetic field pervades all regions of the Sun. In the radiative zone, it is presumably responsible for enforcing the observed solid body rotation \\citep[see][]{MW_87} and hence the dynamical behaviour of the magnetic field is a key factor in the evolution of the radiative zone. For the purposes of this review however, which concentrates on the observed solar magnetic field and the processes responsible for that field, we shall confine our attention to the magnetic field in the convection zone (including the tachocline) and the solar atmosphere. A comprehensive discussion of all aspects of stellar magnetic fields can be found in \\citet{Mestel_99}.  \\subsection{Properties of the solar atmosphere} The visible \\lq surface\\rq\\ of the Sun is the \\textit{photosphere}. The temperature is around $6\\,000$\\,K and the radiation is in the form of white light. The thickness of the photosphere is only 500\\,km; the base of the photosphere is where the optical depth is $\\tau_{500} = 1$ and the top is where the temperature reaches its minimum value.  The most obvious surface features are \\textit{sunspots}; an example is shown in Figure~\\ref{fig:sunspot}. Sunspots have been observed in detail since the invention of the telescope in the 17th century; they were first identified as regions of extremely strong magnetic field by Hale (1908). The typical magnitude of the magnetic field in a sunspot is the order of $3\\,000$\\,G or $0.3$\\,Tesla. The magnetic field manifested in sunspots is generated by dynamo action in the solar interior.  High resolution images of the photosphere, away from sunspots, show evidence of convection through \\textit{granules} and \\textit{supergranules}. Granules are approximately 1000\\,km in horizontal extent and cover the entire surface of the Sun. Hot plasma rises up from the solar interior in the bright centres, spreads out across the surface and then cools and sinks into the interior along dark lanes. The lifetime of an individual granule is about 20 minutes, with the granulation pattern continually changing as newly formed granules push aside old ones. The horizontal plasma flow within a granule can reach speeds of around 7\\,km\\,s$^{-1}$.  Supergranules are much larger, with a size of about 14\\,000\\,km. They are not visible in white light but are revealed through Doppler velocity measurements. Like the granules, which lie inside the supergranules, they cover the entire photosphere and are  continually changing. Individual supergranules last for around 24 hours and have much slower plasma speeds of about 0.5\\,km\\,s$^{-1}$.  Above the photosphere, the \\textit{chromosphere} is an irregular layer with a thickness of around 10\\,000\\,km, in which the temperature rises from $6\\,000$\\,K to about $20\\,000$\\,K. The plasma now emits in the wavelength of Hydrogen alpha ($H\\alpha$); features seen in $H\\alpha$ are thought to outline the local magnetic field lines. At the top of the chromosphere, the temperature rises sharply through a thin layer called the \\textit{transition region}.  The \\textit{corona} is the Sun's outer atmosphere. It is visible from Earth during total eclipses of the Sun as a diffuse white region surrounding the Sun. The temperature in the corona increases dramatically, to the order of 1 million~K, and the wavelength of the plasma emission is in the X-ray and Extreme Ultra-Violet ranges. One major problem in solar physics is to explain why the corona is so hot. The corona displays a variety of features that are entirely controlled by the local, evolving coronal magnetic field, including streamers, plumes, and loops.  \\subsection{Modelling philosophy} To reproduce through numerical simulations all the detailed observations of the evolution of the Sun's magnetic field, obtained from various instruments, is certainly currently impossible since there are simply too many distinct timescales and lengthscales to resolve in numerical simulations. Indeed, even with the continuing increase in computational power and affordability, such direct modelling will remain impossible for the foreseeable future. There are many physical and atomic processes required for a complete modelling of the solar atmosphere. For example, computationally demanding processes include (i) radiative transfer effects in the photosphere and chromosphere and (ii) optically thin radiative losses, thermal conduction and heating in the corona. In addition, in the cool photosphere and low chromosphere, partial ionisation must be considered. All of these processes require knowledge of, for example, the ion abundances, whether the plasma is in local thermodynamic equilibrium or not, and whether the plasma is in ionisation balance.  Controlled numerical experiments allow one to specify which physical effects are to be studied. Hence, the modelling philosophy is to consider a simpler initial state, for example one that is in equilibrium, and then modify it in a controlled manner to see how one effect at a time influences the evolution. Incorporating too many variations all at once obscures the important or dominant physical processes.  \\subsection{Outline} The outline of the paper is as follows. Section~\\ref{sec:MHDequations} presents the MHD equations used in modelling the solar dynamo, the Earth's dynamo and the magnetic field evolution in the solar corona. In Section~\\ref{sec:cycle} the observational evidence for the solar cycle is presented. Next, Section~\\ref{sec:dynamo} describes the current state of modelling the generation of the Sun's magnetic field, both the large-scale magnetic field, on the scale of active regions, and the small-scale field, on the scale of the granulation. Section~\\ref{sec:comparison} compares the dynamo mechanisms of the Sun and the Earth, and looks at the similarities and differences in the computational models of both types of dynamo.  Moving to the atmospheric consequences of the solar dynamo, Section~\\ref{sec:coronalB} considers the nature of the coronal magnetic field, and the manner in which it is influenced by the movement of the photospheric field. Returning to the formation of individual active regions, Section~\\ref{sec:emergence} investigates how the strong magnetic field generated in the solar interior actually emerges through the photosphere and rises into the corona, where it interacts with the pre-existing field. During this process, many of the observational phenomena, whose properties are linked to the solar cycle, are seen to occur in a self-consistent manner. The final section summarises the paper. ",
        "conclusions": "The recent theoretical ideas to explain the solar dynamo, covering both the generation of the large-scale and small-scale magnetic fields, have been presented. Small-scale magnetic dynamo action results in any suitably turbulent flow provided the magnetic Reynolds number is sufficiently high. Purely on theoretical grounds, therefore, it is to be expected that the granules and supergranules, confined close to the surface, will, of themselves, act to generate small-scale magnetic field. This idea is confirmed by observations that show that the amount of small-scale field is independent of the solar cycle, and hence is not simply a by-product of the large-scale dynamo.  Understanding the generation of the Sun's large-scale magnetic field ---  the field of the solar cycle --- remains one of the great challenges of astrophysical MHD. Although most dynamo modelling has been performed within the framework of mean field electrodynamics, there are significant problems in applying this theory to the astrophysically relevant, high $Rm$ regime. Current solar dynamo models fall into one of three broad categories: distributed dynamos, in which dynamo action takes place throughout the convection zone; interface dynamos, in which dynamo action is localised in the tachocline and lower convection zone; and flux transport dynamos, in which the two key components of the dynamo cycle (`$\\alpha$' and `$\\omega$') are separated by the extent of the convection zone but are connected by a large-scale meridional cell. As discussed in Section~\\ref{Subsec:solar_dynamos}, there are problems with all three models, though the idea that the solar magnetic field is generated by a flux transport dynamo does seem the least plausible. Progress with the solar dynamo problem will result from a combination of high resolution numerical simulations, both global and local, focusing on the tachocline for example, together with an improved theoretical understanding of the fundamental aspects of MHD turbulence.  The eleven year sunspot cycle has a profound effect on the evolution of the global coronal magnetic field. For example, the amount of open magnetic flux tends to peak two years after solar maximum. This is due to the non-potential nature of the coronal magnetic field that inflates the corona, causing the opening of closed field lines. Properties of solar prominences can be difficult to predict but, by following the motions of the observed surface magnetic fields and, hence, the build-up of stresses in the coronal magnetic field, it is now possible to predict where prominences will form and which are likely to erupt as a CME. The ideas behind the global coronal field models, namely the inclusion of new emerging magnetic fields, the stressing due to shearing by differential rotation, the transport of magnetic flux to the poles, the break up of strong field regions and the decay of active regions through diffusion, can be used to follow the evolution of an individual active region. The time resolution of the order of 30 seconds, full disk observations from the Solar Dynamics Observatory now mean that the line-of-sight magnetic field and the observed photospheric flow patterns can be used as input for this model and the subsequent evolution of the coronal magnetic field determined. Initially, the models will require the real-time surface flow velocity as input and the evolution of the resulting coronal magnetic field will be compared with observations. Next, more generic flow fields can be used to see if a similar evolution is generated and if similar dynamic features are seen. Then, in the future development of this model, it will be possible to predict whether a newly emerging sunspot or active region is likely to become unstable and generate a large solar flare or not.  One by-product of the large-scale dynamo is sunspots. They emerge, evolve and disperse on the order of a few weeks and the interaction between the emergence of these strong magnetic fields and the pre-existing coronal magnetic field creates a variety of dynamic solar phenomena. At present, numerical simulations can identify the physical processes responsible for the emergence. However, current computational resources, mainly computer memory, mean that only the uppermost part of the solar convection zone and the low corona can be modelled. In addition, the simulation timescales are based on resolving the shortest timescales, determined by the propagation of Alfv\\'en waves in the solar corona, and these are significantly shorter than the observed times for flux emergence. Hence, to save on computing time, the magnetic buoyancy of the emerging flux tubes is chosen to produce a faster evolution. Thus, the magnetic field strength at the photosphere increases much faster than would be liked and this means that the emergence process is also faster. Now this mismatch in timescales is not a major issue since the correct ordering of the timescales for each physical process is preserved, i.e.\\ Alfv\\'en wave timescale is shortest and so on. Nonetheless, the issue of timescales will be addressed with the next generation of parallel computers.  The simulations of flux emergence are now being performed with more realistic pre-existing magnetic environments. For example, as discussed in Section~\\ref{subsec:rope}, the formation of a new flux rope, during the simulations of the emergence of magnetic fields, often results in the eruption of dense plasma from the lower corona. How closely are these eruptions related to the observed CMEs? Again the timescales are wrong but it is possible that the basic physical processes involving the removal of the overlying magnetic field (see Section~\\ref{subsec:erupt}) are correct.  Another coronal feature frequently observed is the sigmoid, observed in X-rays. Sigmoids naturally arise in simulations from the emergence of a twisted flux rope. However, some of them seem to occur during the decaying phase of an active region. Nonetheless, although in this case new magnetic flux is not emerging, there does seem to be a shearing motion along the magnetic polarity inversion line and a general increase in the magnetic helicity. Both properties occur naturally in the emergence simulations. So again, the key physical processes are correctly described even when sigmoids do not exhibit an increase in photospheric magnetic flux.  We have described some of the wide variety of problems currently of importance in the generation and evolution of solar magnetic fields. The nature of the generation of the large-scale solar magnetic field is still far from understood, and remains one of the great challenges of astrophysical MHD. The continual improvement in the quality of the observations of the solar surface and atmosphere is helping to improve modelling of the photospheric and coronal fields; with the advances in computational resources, the simulations are now allowing us to compare theory and observations at a level never before achievable."
    },
    "1107/1107.1625_arXiv.txt": {
        "abstract": "Fluorine  ($^{19}$F)  abundances {\\bf (or upper limits)} are derived  in  {\\bf six} extragalactic  AGB carbon stars from the  HF(1-0) R9 line at 2.3358  $\\mu$m in  high resolution  spectra. The  stars belong  to the Local Group  galaxies LMC, SMC  and Carina dwarf spheroidal, spanning more  than a  factor 50  in metallicity.  This is  the first  study to probe  the  behaviour   of  F  with  metallicity  in intrinsic extragalactic C-rich AGB stars.  Fluorine could be measured  only in four of the  target stars, showing a wide range in F-enhancements. Our F abundance measurements together with those recently derived in Galactic AGB carbon stars show a correlation with the observed carbon and $s-$element enhancements. The observed correlations however, display a different dependence on the stellar metallicity with respect to theoretical predictions in low mass, low metallicity AGB models.  We briefly discuss the  possible reasons for this discrepancy.  If our findings are confirmed in a larger number of metal-poor AGBs, the issue  of F production in AGB stars will need to be revisited. ",
        "introduction": "Comparisons between abundances of some  key  elements determined in stars  in different  evolutionary stages  with nucleosynthetic  stellar   models  is  a  fundamental   tool  for  our understanding of stellar interiors. One element that  can add new insight into this is fluorine, since this element is very sensitive to nuclear  reactions involving proton  and/or alpha  captures. Actually, the  origin of  this light  element is  still uncertain  although much observational progress has  been made  in recent  years  in cool  K, M  and Asymptotic Giant Branch  (AGB) stars (e.g. Jorissen  et al.  1992, hereafter JSL; Cunha et  al. 2003-08; Smith et al.  2005; Uttenthaler et al. 2008),  and in post-AGB stars and planetary nebulae (Werner et al.  2005; Otsuka  et al. 2008). From  these observational  studies two  conclusions are derived:  i) AGB  stars constitute  the only evidence  of stellar  F production and, ii) the inferred evolution of the F abundance in the Galaxy, when compared with chemical evolution  models (Renda et al. 2004), requires the  contribution   of  the   additional  sources  proposed   so  far: gravitational supernovae  (Woosley et  al. 1990) and  Wolf-Rayet stars (Palacios et al. 2005).  Evidence  of F  production in  AGB  stars was  found by  JSL two decades ago: F enhancements up to  a factor 30 solar and a correlation of F with the  C/O ratio was first reported in Galactic (O-rich \\& C-rich) AGB  stars of solar  metallicity. Since the  C/O ratio is expected to  increase in the  envelope during the AGB phase as a  consequence of the third dredge-up (TDU) episodes, this    was   interpreted    as   evidence    of   F production. The F enhancements found by JSL, however, could not  be accounted  by  detailed models  for  AGB stars  at that  epoch (Mowlavi et al. 1998), or  by more recent ones (Lugaro  et  al. 2004)  that  included  the impact  of the  variation of some nuclear  reactions with uncertain rates affecting F. Nonetheless, a recent re-analysis of the JSL's stars by Abia et al. (2009-10; hereafter Paper I and II) using improved line lists and model atmospheres found that the F  abundances reported in JSL have been overestimated because of  a  possible  lack  of  proper accounting  for  C-bearing  molecule blends. This reanalysis reported F abundances by  $\\sim 0.7$ dex lower on  average.  The  new F  abundances  and  the  observed correlation  between  F and  $s-$element  enhancements  in solar metallicity C-stars, are now fully accounted by the current nucleosynthetic modelling  of low mass AGB  stars (Cristallo  et  al. 2009-11).  At   low metallicities the    situation   is   less clear.  Theoretically, the  F production in  AGB  stars is expected to increase when decreasing metallicity.  AGB stars synthesise  F  via $^{14}$N$(n,p)^{14}$C~$(\\alpha,\\gamma)^{18}$O$(p,\\alpha)^{15}$N$(\\alpha,\\gamma)^{19}$F and $   $ $^{14}$N$(\\alpha,\\gamma)^{18}$F$(\\beta+)^{18}$O$(p,\\alpha)^{15}$N$(\\alpha,\\gamma)^{19}$F, where the neutrons are provided by $^{13}$C$(\\alpha,  n)^{16}$O and the protons mainly by $^{14}$N$(n,  p)^{14}$C.  Thus,  the  production  of F basically  depends on  the  availability of  $^{13}$C  in the  He-rich intershell, but also on the amount of $^{13}$C available in the ashes of  the H-burning shell. The  former is weakly  dependent {\\bf on} the stellar metallicity but the  latter scales  with the CNO abundances in the  envelope, which, in turn, depends  on the (primary) $^{12}$C dredged up during TDU episodes. As a consequence, the resulting fluorine can be roughly considered  a primary element. Because of this primary nature, larger  [F/Fe]\\footnote{We  adopt  the usual notation with [x/y]$=$log    $[N(x)/N(y)]_{\\star}-$     log $[N(x)/N(y)]_{\\sun}$    and    log    $\\epsilon(x)=    12    +    $log $[N(x)/N(H)]$.} ratios are obtained in  metal-poor, low-mass  AGB   models  as  compared  with  the  solar metallicity  case.  Nonetheless, the available determinations of F abundances in low metallicity stars depict a more complex scenario than that expected from this simple theoretical {\\bf scheme}. For  instance,  Cunha  et al. (2003) derived [F/O]$\\la 0.0$ in RGB stars of $\\omega$  Cen {\\it enriched} in  $s-$process  elements, which  is difficult to  reconcile with metal-poor AGB stars  being significant F contributors.  Recently, Lucatello  et  al.  (2011) found  lower [F/Fe]  ratios than  theoretically expected  in a  sample  of Galactic carbon enhanced metal-poor stars  enriched in s-elements (CEMP-s). The implications  of  these findings  are  not  easy: first, the surface composition in the  $\\omega$ Cen giants is  hampered by  the  uncertain star  formation history and  chemical evolution  of  this  cluster  and,  on  the  other  hand, the abundances measured in CEMP-s stars (mostly binaries) {\\bf might be affected by the} dilution  of the  accreted material onto the envelope of the secondary star. The amount of material accreted  is  uncertain and, thus are {\\bf in general} the interpretation  of  the  abundances.  On  the  contrary,  F abundances derived in intrinsic metal-poor AGB stars, i.e. stars which own   their  chemical   peculiarities   to  internal   nucleosynthesis, are free of these  problems\\footnote{In Section 3 we show that the [F/Fe] ratio achieved in the envelope in the C-rich AGB phase is almost independent on the initial F content in the star.}. Unfortunately, the few metal poor Galactic AGB stars, as those observed in globular clusters, have masses too low to undergo TDU events and thus, to  pollute the envelope  with the  nucleosynthesis  ashes of  the He-rich  intershell. For  this reason,  the  AGB C-stars in  the satellite galaxies  of the  Milky Way offer  a unique  opportunity to study  the  $^{19}$F  production  at low  metallicity.  These  galaxies are well  known to content metal-poor stellar populations.  Here  we  present for  the  first  time  F abundance  measurements  in metal-poor AGB  carbon stars in stellar systems other  than the Galaxy: the SMC, LMC and Carina dSph galaxies. These C-stars are  significantly  more metal  poor  than the  Galactic counterparts  so  far  analysed  providing  valuable information on  how  F  is produced in AGB stars as a function of metallicity. ",
        "conclusions": "Table 1 summarises the main results. We find a wide range in [F/Fe] ratios in metal-poor extragalactic AGB C-stars. Fluorine is enhanced in all except one star, confirming that this element is produced during the AGB phase. Figure 2 provides further evidence of this. The similarity in the run of F with C strongly support that the synthesis of F is tied to the production of C during the Thermal Pulsing (TP) AGB phase. Also Figure 2 qualitatively shows that F enhancement increases for decreasing stellar metallicity (although the star B30 seems to deviate from this trend). In Paper II we showed that at solar metallicity, the F and C correlation can be reproduced by current TP-AGB models (Cristallo et al. 2009, continuous lines in Figure 2). Models and observations, however, now seem to disagree in metal-poor AGB stars (dashed lines); theoretical models tend to overproduce C for a given F enhancement. Indeed, the two stars in Carina should be along the model computed with $Z\\sim 10^{-4}$ (right dashed line), while the star B30 along the model line with $Z\\sim 10^{-3}$. This discrepancy, which has been already noted at solar Z (e.g. Abia et al. 2002), becomes more evident at lower metallicity as Figure 2 shows, since in current AGB models the efficiency of the TDU increases at low metallicity and thus, the amount of C dredged-up into the envelope. Note that this problem could be also an observational bias: intrinsic AGB C-stars with high C/O ratios should exist, as high values of C/O have been observed both in post-AGB stars and in extrinsic C-rich objects originated by mass transfer from an AGB star. However, in cool C-stars an excess of carbon is immediately translated into a copious production of C-rich dust and into a high opacity of the circumstellar environment, to the point of hiding the photosphere at visual wavelengths. Most of the C might be trapped into grains. These dusty C-stars should show large infrared excess. Curiously enough, the two stars in our sample (Table 1) with high C/O ratios do not show this as deduced from their available infrared photometry, while the star GM780, with large infrared excess (see Section 2) has the typical C/O$\\ga 1$ ratio observed in C-stars.   Figure 3 provides another piece of information on the synthesis of F in AGB stars. It shows the observed relation between F and the average\\footnote{Obtained as the mean value of the [Sr, Zr, Y, Ba, La, Nd, Sm/Fe] ratios.} $s$-element enhancements in C-stars of different metallicities compared with theoretical predictions. Again the observed trend {\\bf points-out} to a correlation between [F/Fe] and [$<$s$>$/Fe]. Despite the scarce number of metal-poor objects studied, the observations suggest a different dependence than that theoretically expected (Cristallo et al. 2011) as a function of the metallicity: in metal-poor C-stars larger [F/Fe] ratios are expected for the observed $s$-element enhancement. While the F and Na enhancements\\footnote{Na can be also produced during the TP-AGB phase, see Cristallo et al. (2009).} found in ALW-C7 can be simultaneously accounted (within error bars) by our reference 1.5 M$_\\odot$, $Z\\sim 10^{-4}$ TP-AGB model, the predicted [$<$s$>$/Fe] ratio is lower than the observed value. It is a challenge that the same models which nicely reproduce the trend found at solar metallicity (lower dashed line in Figure 3), fail in reproducing the metal-poor data. This discrepancy can also be noticed in Figure 4, where the [F/$<$s$>$] ratios found in Galactic CEMP-s (Lucatello et al. 2011) and intrinsic AGB C-stars are plotted. Note that by plotting the [F/$<$s$>$] ratio we avoid any dilution problem that may affect the abundances derived in CEMP-s stars. From this figure, it is clear again that the F content derived in metal-poor stars is lower than the predicted values.  Now the question is how to account for these apparent low F enhancements? In the recent years, some degree of extramixing processes have been invoked to explain the $^{16}$O/$^{17}$O/$^{18}$O ratios found in grains of AGB origin, and the low $^{12}$C/$^{13}$C ratios in C-stars (Busso et al. 2010). Extramixing has been proposed by Denissenkov et al. (2006) {\\bf as a possible solution}{\\footnote{{\\bf There are other explanations more widely accepted, see e.g. Marino et al. (2008).}}} of the observed anti-correlation of F and Na abundances found in low-mass giants of M4 (Smith et al. 2005). Extramixing might reduce F in the AGB envelope if such processes expose material at temperatures high enough to activate the $^{19}$F$(p,\\alpha)$ reaction. However, the $^{16}$O/$^{17}$O ratios found in TRM88 and B30 are those characteristic of stars which have undergone the first dredge-up. The measurement of the $^{16}$O/$^{18}$O ratio in these stars would be critical to give a definite answer on the ocurrence or not of extramixing processes but, unfortunately, $^{18}$O could not be determined. Even the large ($>100$) $^{12}$C/$^{13}$C ratio found in B30 (de Laverny et al. 2006) cannot add any hint on the presence of extramixing, due to the larger $^{12}$C/$^{13}$C ratios attained on the surface according to metal-poor AGB models already after the first TDUs. On the other hand, several nuclear reaction rates affecting F are still uncertain, in particular $^{15}$N$(\\alpha,\\gamma)^{12}$C. We performed {\\bf preliminary} tests with nuclear rates modified within the current uncertainties and found that this can only mildly improve the situation but actually it seems not enough to solve the problem. Another possibility is that these metal-poor C-stars formed with an initial F abundance much lower than the one scaled to their metallicity (see Section 1). Note that the chemical evolution of these satellite galaxies is essentially unknown (even the run of F with [Fe/H] in our Galaxy is uncertain). Nonetheless, according to our theoretical metal-poor AGB models the F production is so large that the [F/Fe] ratio reached in the envelope in the C-rich phase (C/O$>1$) is independent of the initial F content in the star. For instance, in a 1.5 M$_\\odot$, Z$\\sim 10^{-3}$ AGB model assuming an initial [F/Fe]$=-1.0$, the predicted [F/$<$s$>$] ratio after a few TPs {\\bf differs less than 0.1 dex with respect to that obtained with a solar scaled initial ratio. The difference is even smaller for lower metallicity models.}  In summary, from the comparison between the derived F abundances and the current models we conclude that the F synthesis in metal-poor AGB stars is probably not {\\bf as large as expected or some physical mechanism, not currently considered in the models, efficiently destroys it}. Obviously, additional F abundance measurements in metal-poor AGB stars would be extremely important to enlighten this problem. The origin of this element still remains unknown."
    },
    "1107/1107.3620_arXiv.txt": {
        "abstract": "It has been suggested that the high metallicity generally observed in active galactic nuclei (AGNs) and quasars originates from ongoing star formation in the self-gravitating part of accretion disks around the supermassive black holes. We designate this region as the star forming (SF) disk, in which metals are produced from supernova explosions (SNexp) while at the same time inflows are driven by SNexp-excited turbulent viscosity to accrete onto the SMBHs. In this paper, an equation of metallicity governed by SNexp and radial advection is established to describe the metal distribution and evolution in the SF disk. We find that the metal abundance is enriched at different rates at different positions in the disk, and that a metallicity gradient is set up that evolves for steady-state AGNs. Metallicity as an integrated physical parameter can be used as a probe of the SF disk age during one episode of SMBH activity. In the SF disk, evaporation of molecular clouds heated by SNexp blast waves unavoidably forms hot gas. This heating is eventually balanced by the cooling of the hot gas, but we show that the hot gas will escape from the SF disk before being cooled, and diffuse into the BLRs forming with a typical rate of $\\sim 1\\sunmyr$. The diffusion of hot gas from a SF disk depends on ongoing star formation, leading to the metallicity gradients in BLR observed in AGNs. We discuss this and other observable consequences of this scenario. ",
        "introduction": "It is widely accepted that active galactic nuclei (AGNs) and quasars are powered by accretion onto supermassive black holes (SMBHs) (Lynden-Bell 1969). The standard accretion disk model (Shakura \\& Sunyaev 1973) or its variants that account for the so-called \"big blue bump\" (Shields 1978; Malkan 1983; Laor \\& Netzer 1989; Sun \\& Malkan 1989; Brunner et al. 1997; Brocksopp et al. 2006; L\\\"u 2008) show that a typical accretion rate is a fraction of the Eddington limit. Over the last four decades, much attention has been paid to the radiation process and dynamics of the accretion disk, but far less is known about the mechanism driving such an accretion rate. There is little understanding of how the SMBHs are fed at the $\\sim 1$ pc scale, or about the outer boundary condition of the Shakura-Sunyaev accretion disk.  Two important clues can help resolve these wide-open issues. The first is the role of the vertical self-gravity of the outer part of the accretion disk. Paczynski (1978) pointed out that quasar accretion disks are so massive that self-gravity dominates. As a consequence, the disk unavoidably collapses into clumps, among which collisions provide a natural mechanism for viscosity. This was recently stressed again by Rafikov (2009). However, Kolykhalov \\& Sunyaev (1980) suggested that star formation cannot be avoided in such a massive disk. Since this pioneering work, many attempts have been made to construct models that incorporate star formation in the disk as well as take account of its roles in the viscosity and related issues (Sakimoto \\& Coroniti 1981; Lin \\& Pringle 1987; Shlosman \\& Begelman 1987; Shlosman et al. 1989; Shlosman \\& Begelman 1989; Gammie 2001; Goodman 2003; Collin \\& Zahn 1999, 2008; Thompson et al. 2005; Tan \\& Blackmann 2005; Nayakshin \\& Sunyaev 2005; Nayakshin et al. 2008; and the recent numerical simulations of Hobbs et al. 2011; Jiang \\& Goodman 2011). All of these",
        "conclusions": "In this paper, we have extended the study by Wang et al. (2010) of the properties of star forming disks within the central 1 pc regions of quasars and AGNs, where star formation can drive an adequate fuel supply to the SMBH through SNexp. We have made progress in the following areas:  \\begin{itemize} \\item A metallicity equation is derived, unveiling a metallicity gradient in the star forming disk. We find that metallicity can reach up to $\\sim 10-20Z_{\\odot}$ in the innermost regions of the disk and a few times solar abundance in the outer part. As an integrated physical parameter, metallicity is evolving even in disks that dynamically are in a steady state, and the metallicity can reach saturation in a few $10^7$years. Metallicity can be used as a \"clock\" for indicating the age of a star forming disk during one episode of SMBH activity.  \\item Thermal diffusion with a mass rate of $\\sim 1\\sunmyr$ inevitably developes from the star forming regions. This rate is comparable to the outflow rates estimated from X-ray spectroscopic observations. The hot gas diffused above the disk is likely to be the source of the observed broad-line-region gas.  \\item The BLR should then share the metallicity gradient of the SF disk. The metallicity gradient might already have been measured by the systematic difference in the metallicity measured by the low (e.g. \\niiiciii) and high (\\nvciv) ionzation lines in AGNs and quasars. \\end{itemize}  The properties of the star forming disk have important theoretical and observational consequences. It is well known that AGNs have many separate ingredients, including the BLR (with separate HIL and LIL regions), the NLR, the accretion disk, the self-gravitating disk and the torus. However, the origins of these different components as well as the physical connections among them are extremely poorly understood. Since the diffusion rates of the hot gas evaporated from the SF disk are comparable with the accretion rates of the SMBHs, the presence of this gas will have significant consequences. The ionized gas exposed to the AGN irradiation will have a very complicated fate in light of AGN heating, cooling and coupling with its dynamics. However, we speculate that the accumulation of the diffused gas continually supplied from the SF disk can drive the formation of BLRs through dynamical and thermal instabilities. Dependence of the properties of the diffused hot gas on the radius would then lead to onion-like structures within the broad line regions.  It is often suggested that broad line regions, broad absorption lines, and ``warmer\" absorbers are the result of outflows (e.g. Murray \\& Chiang 1997). The present scenario provides a gas supply to the broad line regions, warmer absorbers etc. This points towards a unified model in which broad line clouds, outflowing clouds and warmer absorbers are intrinsically linked with the star forming disk, and hence with the accretion disk around the SMBHs. Such a model will be presented in a forthcoming paper (Wang et al. 2011 in preparation).  A direct proof of the presence of massive stars in the core of a SF disk would be the detection of the spectral features of red supergiants. In the near future, ALMA (Atacama Large Millimeter Array) spectra of galactic nuclei should have the sensitivity to detect the CO bandhead in the near IR that is predicted by our scenario. Another approach would be to use powerful instruments to search for Ca {\\sc ii} triplet ($\\lambda\\lambda 8498, 8542, 8662$\\AA) absorption from the nuclear regions.  Finally, newly born AGNs and quasars may have rapidly growing accretion rates with self-similar behaviors such as $\\siggas\\propto R^{\\gamma_1}t^{\\gamma_2}$ (Yu et al. 2005, but see Schawinski et al. 2010). This would allow us to obtain the general solution for the metallicity evolution during the early stage of the AGN phase. Interestingly, the metal-poor quasars have metallicities similar to their host galaxies ($Z\\le 0.6Z_\\odot$), but they are only a small fraction of quasars (Fu \\& Stockton 2007). As has been suggested by Silk \\& Rees (1998) and others, Ultra-Luminous Infrared Galaxies (ULIRGs) with higher star formation rates (e.g. Laag et al. 2010) and low metallicity are likely to be in the early stage of AGN activity, soon after their birth. Our work here helps move us into a position to apply much more quantitative tests to this general idea."
    },
    "1107/1107.6005_arXiv.txt": {
        "abstract": "Close-in, massive exoplanets raise significant tides in their stellar hosts. We compute the radial velocity (RV) signal due to this fluid motion in the equilibrium tide approximation. The predicted radial velocities in the observed sample of exoplanets exceed $\\rm 1\\ m/s$ for 17 systems, with the largest predicted signal being $\\rm \\sim 30\\ m\\ s^{-1}$ for WASP-18 b. Tidally-induced RV's are thus detectable with present methods. Both tidal fluid flow and the epicyclic motion of a slightly eccentric orbit produce an RV signal at twice the orbital frequency. If care is not taken, the tidally induced RV may, in some cases, be confused with a finite orbital eccentricity. Indeed, WASP-18 b is reported to have an eccentric orbit with small $e=0.009$ and pericenter longitude $\\omega=-\\pi/2$. Whereas such a close alignment of the orbit and line of sight to the observer requires fine tuning, this phase in the RV signal is naturally explained by the tidal velocity signature of an $e=0$ orbit. Additionally, the equilibrium tide estimate for the amplitude is in rough agreement with the data. Thus the reported eccentricity for WASP-18 b is instead likely a signature of the tidally-induced RV in the stellar host. Measurement of both the orbital and tidal velocity for non-transiting planets may allow planet mass and inclination to be separately determined solely from radial velocity data.  We suggest that high precision fitting of RV data should include the tidal velocity signal in those cases where it may affect the determination of orbital parameters. ",
        "introduction": "Increasing RV precision allows the detection of smaller and/or more distant planets. The current state of the art is Doppler measurement errors in the range $ \\sim \\rm 1 \\ m\\ s^{-1}$ (e.g. \\citealt{1996PASP..108..500B,2006SPIE.6269E..23L}). At this level of measurement precision, other physical effects become potentially detectable besides the Keplerian stellar orbital motion: the Rossiter-McLaughlin effect (e.g. \\citealt{2011EPJWC..1105002W}), additional planets perturbing the orbit (e.g. \\citealt{2001ApJ...551L.109L}), ``jitter\" due to convective overturn motions in the stellar envelope \\citep{2005PASP..117..657W} and solar-like acoustic waves \\citep{2007CoAst.150..106B}.  In this paper we investigate an additional effect, the time-dependent spectroscopic shift due to fluid motions in the star forced by the planetary tide.  Previous investigations have highlighted that tides raised in the star by the planet may give observable ellipsoidal (flux) variation \\citep{2003ApJ...592.1217S, 2003ApJ...588L.117L,2008ApJ...679..783P}. NASA's Kepler Mission has recently announced the first detection of ellipsoidal variation due to a planet \\citep{2010ApJ...713L.145W}.  The tidal velocity signal has been discussed by \\citet{2002A&A...384..441W} for the case of a massive star primary and stellar mass compact object secondary. \\citet{1977AcA....27..203D} gives formulae to convert the surface fluid motions to disk averaged radial velocities along the line of sight to the observer. In this paper we discuss the tidally induced fluid motions due to a planetary companion with short orbital period, for which the tidal velocity may be observable.  The tidal velocity is computed in \\S \\ref{sec:rv}. Distinguishing between a small orbital eccentricity and the tidal velocity is discussed in \\S \\ref{sec:distinguish}. Equilibrium tide estimates for the tidal velocity of known exoplanets, and the cases of WASP-18 b and WASP-33 b, are discussed in \\S \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\begin{table*} \\caption{ Estimates of Tidal Velocity } \\label{tab:estimates} \\begin{tabular}{lccccccccc} \\hline planet name &  $M_p[\\sin i]/M_J$ & $P_{\\rm orb}/\\rm day$ & $e$ & $\\omega[deg]$ & $i[deg]$ & $M/M_\\odot$  & $R/R_\\odot$ & $K_{\\rm tide}[m/s]^{\\rm a} $ & $e K_{\\rm orb}[m/s] $ \\\\ \\hline WASP-19 b                         &       1.17   &       0.79   &     0.0046   &       3.00   &      79.40   &       0.97   &       0.99   &       2.83   &       1.19  \\\\ WASP-43 b                         &       1.78   &       0.81   &              &              &      82.60   &       0.58   &       0.93   &       8.90   &             \\\\ WASP-18 b                         &      10.11   &       0.94   &     0.0085   &     -92.10   &      86.63   &       1.24   &       1.36   &      31.94   &      15.38  \\\\ WASP-12 b                         &       1.40   &       1.09   &              &              &      86.00   &       1.35   &       1.60   &       4.78   &             \\\\ OGLE-TR-56 b                      &       1.30   &       1.21   &              &              &      78.80   &       1.17   &       1.32   &       2.12   &             \\\\ HAT-P-23 b                        &       2.09   &       1.21   &     0.1060   &     118.00   &      85.10   &       1.13   &       1.20   &       3.61   &      39.03  \\\\ WASP-33 b                         &       4.59   &       1.22   &     0.174$^{\\rm c}$ & -89.0$^{\\rm c}$ &      87.67   &       1.50   &       1.44   &       5.89   & 120.0            \\\\ HD 41004 B b                      &      18.40   &       1.33   &     0.0810   &     178.50   & no transit   &       0.40   &       0.48$^{\\rm b}$   &       3.61   &     517.48  \\\\ WASP-4 b                          &       1.12   &       1.34   &              &              &      88.80   &       0.93   &       1.15   &       1.13   &              \\\\ CoRoT-14 b                        &       7.60   &       1.51   &              &              &      79.60   &       1.13   &       1.21   &       4.11   &              \\\\ SWEEPS-11                         &       9.70   &       1.80   &              &              &      84.00   &       1.10   &       1.45   &       7.04   &              \\\\ HAT-P-7 b                         &       1.80   &       2.20   &              &              &      84.10   &       1.47   &       1.84   &       1.02   &              \\\\ WASP-14 b                         &       7.72   &       2.24   &     0.0903   &     254.90   &      84.79   &       1.32   &       1.30   &       1.55   &      90.59  \\\\ OGLE2-TR-L9 b                     &       4.34   &       2.49   &              &              &      79.80   &       1.52   &       1.53   &       1.85   &             \\\\ XO-3 b                            &      11.79   &       3.19   &     0.2600   &     345.80   &      84.20   &       1.21   &       1.38   &       3.07   &     384.77  \\\\ HAT-P-2 b                         &       8.74   &       5.63   &     0.5171   &     185.22   &      90.00   &       1.36   &       1.64   &       5.04   &     482.55  \\\\ HIP 13044 b                       &       1.25   &      16.20   &     0.2500   &     219.00   & no transit   &       0.80   &       6.70   &       2.74   &      30.01  \\\\ \\hline \\end{tabular} \\\\ $^{\\rm a}$ Semi-amplitude ((max-min)/2) found by evaluating eq.\\ref{eq:vtide} over the orbit. \\\\ $^{\\rm b}$ Radius of HD 41004 B b  not listed in exoplanets.eu, so we use the main sequence value of $0.48\\ R_\\odot$. \\\\ $^{\\rm c}$ The values of $e$ and $\\omega$ for WASP-33 b were found by \\citet{2011MNRAS.tmp.1075S}, but rejected as unphysical. Their main results instead fix $e=0$. \\\\ \\end{table*}   To evaluate the magnitude of the tidally-induced radial velocity (eq.~[\\ref{eq:vtide}]) for observed exoplanets, data were taken from the Extrasolar Planet Encyclopedia (http://exoplanet.eu/catalog.php) with two exceptions. Updated values for WASP-18 b were taking from \\citet{2010A&A...524A..25T}, and for WASP-33 b we include the eccentricity fit from \\citet{2011MNRAS.tmp.1075S}, and use the upper limit on the mass as an estimate of the mass itself. For transiting planets, the measured inclination is used. For non-transiting planets, we set $\\theta_o=\\pi/2$ and use the measured $M_p\\sin i$ in place of $M_p$.  Parameters used and the tidal velocity semi-amplitude, $K_{\\rm tide}$, found by evaluating equation (\\ref{eq:vtide}) over an orbit, are presented in Table \\ref{tab:estimates} for planets with $K_{\\rm tide}>1\\ {\\rm m\\ s^{-1}}$. For the eccentric orbit cases, the $k=2$ orbital velocity amplitude $e K_{\\rm orb}$ is given to compare to $K_{\\rm tide}$.  Most cases show massive, close-in planets in nearly circular orbits around stars with radii slightly larger than the Sun. Some planets, such as HAT-P-2 b, XO-3 b and HIP 13044 b have enhanced signal near pericenter due to large eccentricity.     Table \\ref{tab:estimates} shows that the tidal velocity is potentially detectable in up to 17 known exoplanets for an RV precision of $1\\ {\\rm m\\ s^{-1}}$.  Inspection of the RV curves in the NsTED database (http://nsted.ipac.caltech.edu/) shows that in many cases, the precision of the measurements taken was not high enough to allow measurement of signals in the range $v_{\\rm tide}=1-30\\ {\\rm m\\ s^{-1}}$. Should such measurements become available in the future, measurement of the tidal velocity will allow a non-trivial check on the planet mass, inclination, and stellar mass and radius. For non-transiting planets detected only through RV data, the tidal signal and orbital motion will allow $M_p$ and $i$ to be separately determined solely from RV data, since they have different dependence on inclination. As signal to noise can be built up over many orbits, follow-up observations on planets already detected may allow one to measure planet masses for a number of RV and transiting planets.  We suggest that a model of the tidal RV should be included in high precision analyses of RV data for close-in planets such as those in Table \\ref{tab:estimates}. At present, the parameter $f_2$ in the tidal RV (eq.~[\\ref{eq:vtidecirc}]) -- which characterizes the properties of the tide raised by the planet -- could be included as a parameter in the fit. In the future, more detailed calculations of the tidal response of the star to its companion may provide accurate values of $f_2$ for the required range of stellar masses and orbital periods.  We have for simplicity focused on the equilibrium tide predictions ($f_2 \\simeq 1.1$) in this paper.  WASP-18 b clearly stands out as the best candidate for an existing detection of the tidal velocity, as the claimed orientation of the eccentric orbit, $\\omega\\simeq-\\pi/2$ \\citep{2010A&A...524A..25T}, is in excellent agreement with the phase predicted by the tidally-induced RV.  Moreover, the equilibrium tide prediction for the amplitude is within a factor of 2 of the inferred $e K_{\\rm orb}$. Quantitatively, we find that instead of the equilibrium tide value, $f_2^{\\rm (eq\\ tide)}=1.1$, the value measured from the data for WASP-18 b is $f_2^{\\rm (data)} = 1.1 \\times (15.38/31.94) = 0.53$.  In our preliminary numerical calculations of the full tidal response of a star, we have found that the equilibrium tide approximation is only accurate at the factor of 2 level for calculating the tidal RV signal (particularly for the horizontal motion at the surface).  Thus we regard the agreement between the data and the equilibrium tide prediction of the RV amplitude for WASP-18 b as satisfactory.  Given this agreement, we argue that the simplest interpretation of the data is that the signal detected in WASP-18 b is the tidal velocity, rather than a finite orbital eccentricity, and that the true orbital eccentricity has a value $e \\ll 0.009$.  Due to the large value of $R/a$, higher harmonics of the orbital frequency due to multipole orders $\\ell=3,4$ may be detectable in the data as well.    WASP-33 b is another candidate in which tidal RV may have been mistaken for an eccentric orbit. \\citet{2011MNRAS.tmp.1075S} initially fitted an eccentric orbit, finding $e=0.174$ and $\\omega = - 89^\\circ$. While the value of $\\omega$ is strongly suggestive of tidal velocity, the large amplitude would require fluid motions much larger than the equilibrium tide (see Table \\ref{tab:estimates}). While larger amplitude tidal flow is expected for higher mass stars, due to their thin surface convection zones, this factor of $\\simeq 20$ increase in amplitude likely requires a near resonance with a g-mode.  The case of WASP-33 b brings up the issue of how stellar rotation affects the tidal RV signal. While most planet-host stars have slow, sub-synchronous rotation (even WASP-18 b, an A star) WASP-33 b is an example of an A star that rotates more rapidly than the orbit. Stellar rotation can be included in the analysis in two places, through alteration of the fluid motions, mainly through the Coriolis force, and also by changing the expression for the fluid motion at the perturbed surface of the star. As the equilibrium tide ignores fluid inertia, the equilibrium tide displacements are unchanged for a rotating star. The second way in which rotation changes the radial velocity is that rather than using $\\dot{\\vec{\\xi}}$ in equation (\\ref{eq:avgvpar}), as is appropriate for a non-rotating star, the Lagrangian velocity perturbation \\be \\Delta \\vec{v} & = & \\dot{\\vec{\\xi}} + \\vec{v} \\cdot \\grad \\vec{\\xi} \\label{eq:ldv1} \\\\ & = & \\vec{e}_i \\left( \\frac{\\partial }{\\partial t} + \\Omega \\frac{\\partial }{\\partial \\phi} \\right)\\xi_i + \\vec{\\Omega} \\times \\vec{\\xi} \\label{eq:ldv2} \\ee must be used, which includes the effect of (here uniform) rotation ($\\vec{v} = \\vec{\\Omega} \\times \\vec{x}$) in the background star. The bracketed term in equation (\\ref{eq:ldv2}) is the time derivative in the frame co-rotating with the star. For a star with rotation synchronized to the orbit, this term is zero as the tide is time-independent in the fluid frame. The second term in eq.\\ref{eq:ldv2} represents the rigid rotation of the deformed surface, and is present even for a synchronized star.  We have evaluated the form of equation (\\ref{eq:ldv2}) for the case of a circular orbit with rotation and orbital angular momentum axes aligned. The resulting expression shows that $n \\rightarrow n-\\Omega$ in the time derivative term, and new terms $\\propto \\Omega$ arise from the piece $\\vec{\\Omega} \\times \\vec{\\xi}$. The expected phase of the radial velocity signal should still be $\\omega = - \\pi/2$ for the $\\Omega < n$ case. However, it may be possible for the tidal RV phase to shift by $180^\\circ$ when $\\Omega > n$. A more detailed calculation is necessary to settle this issue. For WASP-18 b, rotation may introduce a small amplitude correction ($\\sim 15\\%$). The amplitude correction for WASP-33 b is expected to be order unity. We plan to include rotation in a future study."
    },
    "1107/1107.3416_arXiv.txt": {
        "abstract": "{We present an optical spectroscopic atlas at intermediate resolution (8--15\\,\\AA) for 123 core-dominated radio-loud active galactic nuclei with relativistic jets, drawn from the  MOJAVE/2cm sample at 15\\,GHz. It is the first time that spectroscopic and photometric parameters for a large sample of such type of AGN are presented. The atlas includes spectral parameters for the emission lines H$\\beta$, [\\ion{O}{3}]$\\lambda$5007, \\ion{Mg}{2}$\\lambda$2798 and/or \\ion{C}{4}$\\lambda$1549 and corresponding data for the continuum, as well as the luminosities and equivalent widths of the \\ion{Fe}{II} UV/optical. It also contains the homogeneous photometric information in the B-band for 242 sources of the sample, with a distribution peak  at $B_J$=18.0 and a magnitude interval of 11.1$\\,\\leq\\,B_J\\,\\leq\\,$23.7.}   \\resumen{Presentamos un atlas espectrosc\\'opico \\'optico de resoluci\\'on intermedia (8--15\\,\\AA) para 123 n\\'ucleos activos de galaxias compactos con la presencia de chorros superlum\\'inicos, tomados de la muestra limitada en densidad de flujo a 15\\,GHz MOJAVE/2cm. Es la primera vez que se presentan los par\\'ametros espectrosc\\'opicos y fotom\\'etricos  para una muestra tan grande de este tipo de AGN. El atlas incluye los par\\'ametros espectrales para las l\\'ineas de emisi\\'on H$\\beta$, [\\ion{O}{3}]$\\lambda$5007, \\ion{Mg}{2}$\\lambda$2798 y/\\'o \\ion{C}{4}$\\lambda$1549, junto con los datos para la emisi\\'on del continuo correspondiente. Se presentan adem\\'as las luminosidades y el ancho equivalente del  \\ion{Fe}{II} UV/\\'optico. Contiene tambi\\'en la informaci\\'on fotom\\'etrica homog\\'enea en la banda B para 242 fuentes de la muestra, con un pico en la distribuci\\'on de $B_J$=18.0 y un intervalo en magnitud de 11.1$\\,\\leq\\,B_J\\,\\leq\\,$23.7.}   \\addkeyword{atlases\\,---\\,galaxies: active} \\addkeyword{galaxies:\\,nuclei} \\addkeyword{quasars:\\,emission lines} \\addkeyword{techniques:\\,spectroscopic} ",
        "introduction": "In the current paradigm of active galactic nuclei (AGN), the large amount of energy released by the AGN is generated in the very small region, the central engine, which is thought to be powered by accretion of matter onto a central black hole. This active central engine also ejects bipolar, highly collimated, relativistic outflows (jets). AGN in which one of the jets is oriented towards the line of sight is strongly beamed due to relativistic effects. This type of sources poses several fundamental questions such as: how the central engine is related to the pc-scale jets; what is the contribution of the jet emission to the total continuum emission and line emission; how are related the continuum emission region, broad/narrow emission-line region with different properties of the relativistic jet. Despite all the advances in AGN research, a global analysis of the physics involved at all spatial scales, from sub-parsec to kpc, is needed to tackle above issues.  Recently, \\citet{arshakian08,arshakian09} and \\citet{tavares09}, using the long-term optical spectral and radio (VLBI) monitoring data of radio galaxies 3C 390.3 and 3C 120 found a link between the variable optical emission and the kinematics of the sub-parsec-scale jet. They have been able to localize the region of a variable optical emission in the innermost sub-pc scale region of the jet, and showed that very long-term variations ($\\sim$ 10 yr) of optical continuum emission are correlated with the radio emission from the base of the jet located just above the disk, while the optical long-term variations (1--2 yr) follow the radio flares from the stationary component in the jet with a time delay of about one yr. Using the flux-limited complete sample of core-dominated AGN  \\citep[MOJAVE-1;][]{lister09}, \\citet{arshakian10} found interesting correlations among properties of parsec-scale jets, using the photometric data presented in this atlas. They reported a significant positive correlation between optical nuclear emission and total radio emission at 15 GHz for 99 quasars. Radio emission originates in the unresolved core, at milliarcsecond scales suggesting that both radio and optical emission are beamed and originated in the innermost part of the sub-parsec-scale jet in quasars. For BL Lacs, the optical continuum emission correlates with the radio emission of the jet. These results are confirmed for a larger sample of 233 core-dominated AGN \\citep{torrealba11}.  Furthermore, spectroscopic parameters of the broad- and narrow-line profiles provide a direct clue about the physics, kinematics, and structure of the central engine of AGN. Diverse studies have  shown the existence of an anticorrelation between the prominence of the radio nucleus and the width of broad emission line,  giving a clue on the gas distribution in the broad line region (BLR) in these sources \\citep[H$\\alpha$, H$\\beta$, \\ion{Mg}{2}, and others lines; ][]{hough02,vestergaard00,WB86}. The results of \\citet{rokaki03} confirmed the correlation between the jet viewing angles and the broad emission line equivalent widths of 19 superluminal quasars, which suggests a flattened structure for the line-emitting material.  Motivation of the present paper is to perform a robust statistical analysis and to study in detail the physical link between the diverse emission regions in radio-loud AGN from sub-pc to kpc scales. For this purpose it is necessary to compile a well defined sample of compact AGN with superluminal jets and a collection of good quality spectral parameters  (signal--to--noise $\\sim30$). Here, we gathered the spectroscopic and photometric information for a specific subsample of 250 compact AGN (selected from the MOJAVE sample) for which the jet parameters are well characterized \\citep{kovalev05}.  We present an optical spectroscopic atlas of 123 AGN that represent the $\\sim$50\\% of the total sample, supplemented with the photometric information for 97\\%  of the same sample (242 sources). The contents of the paper are as follows: the sample is described in \\S~\\ref{sec:sample}, followed by the photometric data and calibration for estimating the optical luminosities at 5100 \\AA\\, (\\S~\\ref{sec:L5100});  details of optical spectroscopic observations are presented in \\S~\\ref{sec:observ}, followed by data reduction and calibration procedures in \\S~\\ref{sec:data}; the detailed procedure to extract the principal parameters of the continuum emission and various emission lines (H$\\beta$, [\\ion{O}{3}]$\\lambda$5007, \\ion{Mg}{2}$\\lambda$2798, \\ion{C}{4}$\\lambda$1549, and \\ion{Fe}{II} UV/optical) are described in \\S~\\ref{sec:processing}, followed by the results in \\S~\\ref{sec:results}, and  finally, the appendices contain the general properties of the sample in~\\ref{secap:mojave}, spectral atlas in~\\ref{sec:atlas}, and  emission line parameters and continuum emission in~\\ref{secap:parameters}.  Throughout the paper a flat cosmology model is used with $\\Omega_{m}=0.3$ ($\\Omega_{\\Lambda}+\\Omega_{m}=1$) and $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\begin{tabular}{lcccc} \\toprule  \\multicolumn{1}{c}{ {}} & \\multicolumn{1}{c}{ {}} & \\multicolumn{3}{c}{ {Spectral Region}} \\\\ \\cline{3-5}   \\multicolumn{1}{c}{ {Sample}} & \\multicolumn{1}{c}{ {Spectral Type}} & \\multicolumn{1}{c}{ {H$\\beta$}} & \\multicolumn{1}{c}{ {\\ion{Mg}{2}}} & \\multicolumn{1}{c}{ {\\ion{C}{4}}} \\\\  \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} \\\\ \\midrule  MOJAVE/2cm        & BL    & 6  & \\nodata & \\nodata\\\\ & RG    & 12 & 2       & 1 \\\\ & HPRQ  & 8  & 19      & 9 \\\\ & LPRQ  & 7  & 18      & 6 \\\\ & Q     & 8  & 39      & 19 \\\\ \\hline &    Total   & 41 & 78      & 35 \\\\[0.5ex] \\hline  MOJAVE-1          & BL    & 4  & \\nodata & \\nodata\\\\ & RG    & 7  & 1       & \\nodata \\\\ & HPRQ  & 7  & 16      & 9 \\\\ & LPRQ  & 6  & 14      & 6 \\\\ & Q     & 4  & 15      & 8 \\\\ \\hline & Total   & 28 & 46      & 23 \\\\ \\bottomrule  \\end{tabular} \\end{table*}   \\subsection{Tabular data} The results of measurements for each parameter of the emission lines and continuum emission for different regions are presented in the Tables~\\ref{tabA1:hb1} to \\ref{tabA5:CIV_1} of the Appendix~\\ref{secap:parameters}. The format of these tables is explained in what follows.  Table~\\ref{tabA1:hb1} presents the emission line parameters of the region around H$\\beta$ for 41  AGN. Col. (1) is the source name, Col. (2) is the spectrum reference, Col. (3)--(6) are the parameters of the H$\\beta$ broad component: FWHM, EW, total flux, and luminosity, respectively, Col. (7)--(10) are the parameters of the [\\ion{O}{3}] $\\lambda$5007 line: FWHM, EW, flux and luminosity, respectively. Nine of 41 AGN have only narrow emission lines with FWHM H$\\beta\\lesssim$ 1000 km\\, s$^{-1}$: 0238$-$084, 0316$+$413, 0710$+$439, 0745$+$241, 0831$+$557, 1155$+$251, 1345$+$125, 1642$+$690, and 1957$+$405.  Table~\\ref{tabA2:hb2} presents the parameters obtained for the emission continuum 5100\\,\\AA\\, and the \\ion{Fe}{2} $\\lambda4570$: Col. (1) is the source same, Col. (2)--(3) are the flux density and luminosity at 5100\\,\\AA, respectively, Col. (4) is the spectral index of the local continuum, Col. (5)--(7) are the total flux, luminosity, and EW of the \\ion{Fe}{2} $\\lambda4570$, respectively.  Table~\\ref{tabA3:MgII_1} presents the continuum emission at 3000\\,\\AA, and parameters  of the \\ion{Mg}{2}\\,$\\lambda$2798 emission line for 78  AGN: Col. (1) is the source name, Col. (2) is the spectrum reference, Col. (3)--(6) are the parameters of the \\ion{Mg}{2} broad component: FWHM, EW, total flux, and luminosity, respectively, Col. (7)--(8) are the flux density and luminosity at 3000\\,\\AA\\,, respectively, and Col. (9) is the spectral index of the local continuum at 3000\\,\\AA.  The emission line parameters of the \\ion{Fe}{2} $\\lambda2490$ are presented in Table~\\ref{tabA4:MgII_2}: Col. (1) is the source name, Col. (2)--(4) are the total flux, luminosity, and EW, respectively, and Col. (5) is the comment to  \\ion{Fe}{2} $\\lambda2490$  measurements.   The continuum emission at 1350\\,\\AA\\, and parameters of the \\ion{C}{4}\\,$\\lambda$1549 emission line for 35 AGN are presented in Table~\\ref{tabA5:CIV_1}: Col. (1) is the source name, Col. (2) is the spectrum reference, Col. (3)--(6) are the parameters of the \\ion{C}{4} broad component: FWHM, EW, total flux, and luminosity, respectively, Col. (7)--(8) are the flux density and luminosity at 1350\\,\\AA, and Col. (9) is the spectral index of the local continuum at 1350\\,\\AA.   \\subsection{Descriptive statistics for different spectral regions} \\label{sec:statistics} We have measured  the emission line parameters: FWHM, EW, flux, and luminosity, and we obtained the continuum emission flux density and luminosity using our Spectral Atlas of 123 AGN from the MOJAVE/2cm sample. In Table~\\ref{tab5:average_lines_c} we present the descriptive statistics for each spectral parameter.  Columns are as follows: Col. (1) is the parameter, Col. (2) is the spectral type, Col. (3) is the number of sources, Col. (4) is the average value of the parameter, Col.(5) is the standard deviation of the data, and Col. (6) and (7) are the minimum and maximum values of the parameter, respectively. The numbers in parentheses refer to descriptive statistics for the flux-limited sample MOJAVE-1 (M1).  The statistical parameters of the H$\\beta$ region are presented for both radio galaxies and quasars because the data for each spectral type are scarce (4 Galaxies, 7 HPRQ,  7 LPRQ, and 6 Q). Note that the radio galaxy 3C 390.3 was excluded from the analysis because of large FWHM H$\\beta\\,\\rm (BC)\\sim\\,12,000\\,km\\,s^{-1}$, as well the BL Lacs and the narrow line objects. The \\ion{C}{4} region (9 HPRQ, 6 LPRQ, and 19 Q) was treated similarly. For the \\ion{Mg}{2} region the statistical analysis  was performed for different types of quasars (19 HPRQ, 18 LPRQ, and 39 Q). In this case, two radio galaxies were excluded: 0007$+$106 and 3C 390.3.  The broad component of the emission line \\ion{C}{4} has a larger FWHM \\ion{C}{4}\\,(BC)\\,=\\,6498$\\,\\pm\\,$1515\\, km\\,s$^{-1}$ value, between H$\\beta$ and \\ion{Mg}{2}. A caveat here is that the data of each emission line corresponds to a different AGN, and only few sources have data for more than one emission line, as was mentioned in \\S~\\ref{sec:spectral}.  The continuum luminosity at $\\lambda\\,L_{1350}$ is  larger by a factor of $\\sim5$ and $\\sim30$ than the continuum luminosities  $\\lambda\\,L_{3000}$ and $\\lambda\\,L_{5100}$, respectively.   The FWHM of \\ion{Mg}{2} is quite similar for different types of quasars (HPRQ, LPRQ, and Q). This parameter has an average value of FWHM \\ion{Mg}{2}$\\rm\\,(BC)\\,=\\,$5108$\\,\\pm\\,$946\\,km\\,s$^{-1}$.  In contrast,  the EW of \\ion{Mg}{2}$\\rm\\,(BC)\\,=\\,$42$\\,\\pm\\,$29\\,\\AA\\, shows a small difference between types. This parameter for quasars with no polarimetry information (Q) is smaller by 36\\% and  25\\% than those for HPRQ and LPRQ.  The \\textit{KS}-test shows that the EW of the \\ion{Mg}{2} distributions of the HPRQ and LPRQ are different at a confidence level of 95.1\\% ($P_{KS}=0.049$), also the  distributions of the HPRQ and Q are different at a confidence level of 97.8\\% ($P_{KS}=0.011$).  Finally, the \\textit{KS}-test between the EW \\ion{Fe}{2}\\,$\\lambda$2490 distributions of the HPRQ and LPRQ showed no significant difference  ($P_{KS}=0.314$).  Further analysis related to the emission lines parameters and the properties of the pc-scale jets  will be presented in a forthcoming paper. \\newpage \\begin{scriptsize} \\begin{longtable}{lcccccc}  \\caption{Mean values of parameters for continuum and line emission} \\label{tab5:average_lines_c} \\\\  \\hline \\hline \\\\[-2ex] \\multicolumn{1}{c}{Parameter} & \\multicolumn{1}{c}{Type} & \\multicolumn{1}{c}{Number} & \\multicolumn{1}{c}{Average} & \\multicolumn{1}{c}{$\\sigma$} & \\multicolumn{1}{c}{Min.} & \\multicolumn{1}{c}{Max.} \\\\  \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & \\multicolumn{1}{c}{(3)} & \\multicolumn{1}{c}{(4)} & \\multicolumn{1}{c}{(5)} & \\multicolumn{1}{c}{(6)} & \\multicolumn{1}{c}{(7)} \\\\[0.5ex] \\hline \\endfirsthead \\hline \\hline \\endlastfoot FWHM H$\\beta\\rm\\, (BC)$\\,[km\\,s$^{-1}$] &\tAll\t&\t24\t&\t4055\t&\t1331\t&\t1670\t&\t8600\t\\\\ &\t\t&\t(20)\t&\t(3957)\t&\t(856)\t&\t(2650)\t&\t(5750)\t\\\\ EW H$\\beta\\rm\\, (BC)$\\,[\\AA] &\tAll\t&\t24\t&\t56\t&\t24\t&\t13\t&\t125\t\\\\ &\t\t&\t(20)\t&\t(57)\t&\t(26)\t&\t(13)\t&\t(125)\t\\\\ $L_{\\rm H\\beta\\rm\\, (BC)}$\\,[10$^{42}$ erg\\,s$^{-1}$] &\tAll\t&\t24\t&\t40\t&\t50\t&\t1\t&\t182\t\\\\ &\t\t&\t(20)\t&\t(37)\t&\t(44)\t&\t(1)\t&\t(182)\t\\\\\\cline{1-7}  EW \\ion{Fe}{2}\\,$\\lambda$4570\\, [\\AA] &\tAll\t&\t22\t&\t18\t&\t17\t&\t2\t&\t69\t\\\\ &\t\t&\t(17)\t&\t(19)\t&\t(17)\t&\t(2)\t&\t(69)\t\\\\ $L_{4570}$\\,[10$^{41}$ erg\\,s$^{-1}$] &\tAll\t&\t22\t&\t96\t&\t96\t&\t2\t&\t342\t\\\\ &\t\t&\t(17)\t&\t(98)\t&\t(99)\t&\t(2)\t&\t(342)\t\\\\\\cline{1-7}  FWHM [\\ion{O}{3}]\\,$\\lambda$5007\\,[km\\,s$^{-1}$] &\tAll\t&\t35\t&\t767\t&\t382\t&\t360\t&\t1716\t\\\\ &\t\t&\t(24)\t&\t(731)\t&\t(328)\t&\t(360)\t&\t(1394)\t\\\\ EW [\\ion{O}{3}]\\,$\\lambda$5007\\,[\\AA] &\tAll\t&\t35\t&\t39\t&\t62\t&\t0.2\t&\t380\t\\\\ &\t\t&\t(24)\t&\t(43)\t&\t(75)\t&\t(5)\t&\t(380)\t\\\\ $L_{[\\ion{O}{3}]}$\\,[10$^{42}$\\,erg\\,s$^{-1}$]\t&\tAll\t&\t35\t&\t14\t&\t33\t&\t0.004\t&\t191\t\\\\ &\t\t&\t(24)\t&\t(10)\t&\t(13)\t&\t(0.004)\t&\t(50)\t\\\\\\cline{1-7}  $\\lambda\\,L_{5100}$\\,[10$^{44}$\\,erg\\,s$^{-1}$] &\tAll\t&\t41\t&\t31\t&\t42\t&\t0.01\t&\t181\t\\\\ &\t\t&\t(28)\t&\t(34)\t&\t(43)\t&\t(0.006)\t&\t(181)\t\\\\\\cline{1-7}  FWHM \\ion{Mg}{2}$\\rm\\, (BC)$\\,[km\\,s$^{-1}$] &\tAll\t&\t76\t&\t5108\t&\t946\t&\t2445\t&\t8752\t\\\\ &\t\t&\t(45)\t&\t(5174)\t&\t(663)\t&\t(3735)\t&\t(7495)\t\\\\ &\tHPRQ\t&\t19\t&\t5118\t&\t960\t&\t2565\t&\t7495\t\\\\ &\t\t&\t(16)\t&\t(5272)\t&\t(789)\t&\t(4245)\t&\t(7495)\t\\\\ &\tLPRQ\t&\t18\t&\t5151\t&\t1063\t&\t3735\t&\t8752\t\\\\ &\t\t&\t(14)\t&\t(4991)\t&\t(599)\t&\t(3735)\t&\t(5583)\t\\\\ &\tQ\t&\t39\t&\t5083\t&\t908\t&\t2445\t&\t6849\t\\\\ &\t\t&\t(15)\t&\t(5239)\t&\t(578)\t&\t(4287)\t&\t(6266)\t\\\\\\cline{1-7}  EW \\ion{Mg}{2}$\\rm\\, (BC)$\\,[\\AA] &\tAll\t&\t76\t&\t36\t&\t24\t&\t6\t&\t177\t\\\\ &\t\t&\t(45)\t&\t(33)\t&\t(27)\t&\t(6)\t&\t(177)\t\\\\ &\tHPRQ\t&\t19\t&\t27\t&\t18\t&\t6\t&\t74\t\\\\ &\t\t&\t(16)\t&\t(29)\t&\t(19)\t&\t(6)\t&\t(74)\t\\\\ &\tLPRQ\t&\t18\t&\t31\t&\t12\t&\t15\t&\t53\t\\\\ &\t\t&\t(14)\t&\t(30)\t&\t(11)\t&\t(15)\t&\t(53)\t\\\\ &\tQ\t&\t39\t&\t42\t&\t29\t&\t11\t&\t177\t\\\\ &\t\t&\t(15)\t&\t(40)\t&\t(41)\t&\t(13)\t&\t(177)\t\\\\\\cline{1-7}  $L_{\\ion{Mg}{2} \\rm (BC)}$\\,[10$^{42}$\\,erg\\,s$^{-1}$] &\tAll\t&\t76\t&\t177\t&\t695\t&\t7\t&\t6091\t\\\\ &\t\t&\t(45)\t&\t(229)\t&\t(899)\t&\t(11)\t&\t(6091)\t\\\\ &\tHPRQ\t&\t19\t&\t53\t&\t34\t&\t7\t&\t126\t\\\\ &\t\t&\t(16)\t&\t(50)\t&\t(30)\t&\t(11)\t&\t(126)\t\\\\ &\tLPRQ\t&\t18\t&\t476\t&\t1407\t&\t35\t&\t6091\t\\\\ &\t\t&\t(14)\t&\t(587)\t&\t(1591)\t&\t(47)\t&\t(6091)\t\\\\ &\tQ\t&\t39\t&\t99\t&\t97\t&\t12\t&\t438\t\\\\ &\t\t&\t(15)\t&\t(85)\t&\t(60)\t&\t(17)\t&\t(196)\t\\\\\\cline{1-7}  EW \\ion{Fe}{2}\\,$\\lambda$2490\\,[\\AA] &\tAll\t&\t67\t&\t54\t&\t26\t&\t6\t&\t128\t\\\\ &\t\t&\t(40)\t&\t(50)\t&\t(26)\t&\t(6)\t&\t(128)\t\\\\ &\tHPRQ\t&\t18\t&\t41\t&\t21\t&\t6\t&\t79\t\\\\ &\t\t&\t(15)\t&\t(39)\t&\t(22)\t&\t(6)\t&\t(79)\t\\\\ &\tLPRQ\t&\t16\t&\t47\t&\t17\t&\t18\t&\t74\t\\\\ &\t\t&\t(12)\t&\t(47)\t&\t(17)\t&\t(18)\t&\t(67)\t\\\\ &\tQ\t&\t33\t&\t65\t&\t27\t&\t21\t&\t128\t\\\\ &\t\t&\t(13)\t&\t(64)\t&\t(33)\t&\t(22)\t&\t(128)\t\\\\\\cline{1-7}  $L_{2490}$\\,[10$^{42}$\\,erg\\,s$^{-1}$]\t&\tAll\t&\t67\t&\t1614\t&\t10,772\t&\t31\t&\t88,402\t\\\\ &\t\t&\t(40)\t&\t(2456)\t&\t(13,941)\t&\t(40)\t&\t(88,402)\t\\\\ &\tHPRQ\t&\t18\t&\t137\t&\t101\t&\t40\t&\t451\t\\\\ &\t\t&\t(15)\t&\t(118)\t&\t(68)\t&\t(40)\t&\t(275)\t\\\\ &\tLPRQ\t&\t16\t&\t5863\t&\t22014\t&\t36\t&\t88402\t\\\\ &\t\t&\t(12)\t&\t(7760)\t&\t(25,400)\t&\t(100)\t&\t(88,402)\t\\\\ &\tQ\t&\t33\t&\t359\t&\t535\t&\t31\t&\t2760\t\\\\ &\t\t&\t(13)\t&\t(257)\t&\t(153)\t&\t(46)\t&\t(521)\t\\\\\\cline{1-7}  $\\lambda\\,L_{3000}$\\,[10$^{44}$\\,erg\\,s$^{-1}$] &\tAll\t&\t76\t&\t204\t&\t936\t&\t8\t&\t8205\t\\\\ &\t\t&\t(45)\t&\t(287)\t&\t(1212)\t&\t(8)\t&\t(8205)\t\\\\ &\tHPRQ\t&\t19\t&\t88\t&\t70\t&\t13\t&\t237\t\\\\ &\t\t&\t(16)\t&\t(84)\t&\t(70)\t&\t(13)\t&\t(237)\t\\\\ &\tLPRQ\t&\t18\t&\t592\t&\t1905\t&\t23\t&\t8205\t\\\\ &\t\t&\t(14)\t&\t(735)\t&\t(2155)\t&\t(33)\t&\t(8205)\t\\\\ &\tQ\t&\t39\t&\t82\t&\t90\t&\t8\t&\t386\t\\\\ &\t\t&\t(15)\t&\t(84)\t&\t(73)\t&\t(8)\t&\t(223)\t\\\\\\cline{1-7}  FWHM \\ion{C}{4}$\\rm\\, (BC)$\\,[km\\,s$^{-1}$] &\tAll\t&\t34\t&\t6498\t&\t1515\t&\t2818\t&\t9150\t\\\\ &\t\t&\t(23)\t&\t(6640)\t&\t(1500)\t&\t(3445)\t&\t(9150)\t\\\\ EW \\ion{C}{4}$\\rm\\, (BC)$\\,[\\AA] &\tAll\t&\t34\t&\t29\t&\t17\t&\t10\t&\t87\t\\\\ &\t\t&\t(23)\t&\t(29)\t&\t(17)\t&\t(10)\t&\t(87)\t\\\\ $L_{\\ion{C}{4}\\rm\\, (BC)}$[10$^{42}$\\,erg\\,s$^{-1}$] &\tAll\t&\t34\t&\t817\t&\t2063\t&\t20\t&\t11,345\t\\\\ &\t\t&\t(23)\t&\t(857)\t&\t(2314)\t&\t(27)\t&\t(11,345)\t\\\\ $\\lambda\\,L_{1350}$\\,[10$^{44}$\\,erg\\,s$^{-1}$] &\tAll\t&\t34\t&\t983\t&\t3363\t&\t5\t&\t19,566\t\\\\ &\t\t&\t(23)\t&\t(1199)\t&\t(4043)\t&\t(24)\t&\t(19,566)\t\\\\ \\end{longtable} \\end{scriptsize}"
    },
    "1107/1107.4000_arXiv.txt": {
        "abstract": "OSIRIS  (Optical  System for  Imaging  and  low Resolution  Integrated Spectroscopy)  is the first  light instrument  of the  Gran Telescopio Canarias (GTC).  It provides a flexible and competitive tunable filter (TF).  Since  it is based  on a Fabry-Perot interferometer  working in collimated  beam,  the  TF  transmission  wavelength  depends  on  the position of the target with  respect to the optical axis.  This effect is non-negligible and must be accounted for in the data reduction. Our paper  establishes a  wavelength calibration  for OSIRIS  TF  with the accuracy required  for spectrophotometric measurements  using the full field  of  view  (FOV)  of  the  instrument.   The  variation  of  the transmission wavelength $\\lambda(R)$ across  the FOV is well described by     $\\lambda(R)=\\lambda(0)/\\sqrt{1+\\left(R/f_2\\right)^2}$,    where $\\lambda(0)$ is  the central  wavelength, $R$ represents  the physical distance  from the  optical axis,  and $f_2=185.70\\pm0.17\\,$mm  is the effective  focal  length  of  the  camera lens.   This  new  empirical calibration yields an accuracy better than $1$\\,\\AA\\ across the entire OSIRIS   FOV  ($\\sim$8\\arcmin$\\times$8\\arcmin),   provided   that  the position of the  optical axis is known within  45~$\\mu$m ($\\equiv$ 1.5 binned  pixels).  We suggest  a  calibration  protocol  to grant  such precision over  long periods, upon re-alignment of  OSIRIS optics, and in  different wavelength  ranges.  This  calibration differs  from the calibration in OSIRIS manual  which, nonetheless, provides an accuracy $\\lesssim1$\\AA\\, for $R\\lesssim 2\\arcmin$. ",
        "introduction": "\\label{sec:introduction}  Tunable  filter  (TF)  instruments  are becoming  standard  tools  for mapping  physical properties  in extended  astronomical  sources.  Most  modern   large  telescopes  have  TFs   based  on  Fabry-Perot interferometers (FPI; or etalons) such as the 11\\,m Southern African Large Telescope (SALT; Rangwala  et al.  2008) or the Magellan-Baade 6.5 m  telescope \\citep{veilleux10}.  The tunable  imaging era began with the  instrument of \\citet{athertonreay81} and  many others have further developed  the technique \\citep{blandtully89,brown94}.  The current instrumental  design is mainly  based on the  Taurus Tunable Filter  (Bland-Hawthorn  \\&  Jones  1998) formerly  mounted  at  the Anglo-Australian  Telescope  and  afterward  moved  to  the  William Herschel Telescope at La Palma.   The power of these instruments to perform   spectrophotometric  studies  is   clear.   They   allow  for spectrophotometry with an unlimited  number of wavelengths that can be scanned within a  wavelength range.  These etalon based  TFs provide a central monochromatic field (MF).  Outside  the MF, and due to the use of  etalons  in collimated  beam,  there  is  a significant  shift  of wavelength  that   depends  on  the  distance  to   the  center.   The instruments are designed to provide the largest possible MF, the wider possible  wavelength  range, and  the  narrower possible  transmission band-pass.   Besides,  a proper  calibration  corrects the  wavelength displacement with  radius, thus recovering  much of the  field outside the MF to be used as an extended MF.  Among the large  facilities hosting such instruments is  the 10~m Gran Telescopio  Canarias (GTC),  located  at the  Roque  de los  Muchachos Observatory, that started its operations in the first semester of 2009 with OSIRIS \\citep{cepa05,cepa10} as the first light instrument. OSIRIS stands for  \"Optical System for Imaging and  low Resolution Integrated Spectroscopy\".   Besides  conventional  imaging and  spectroscopy,  it provides the additional  capability of a TF observing  mode working in the visible.  It  covers the wavelength range 0.365-1.0  $\\mu$m with a nominal  FOV of  7.8\\arcmin$\\times8.5$\\arcmin.   Further commissioning tests     establish      the     real     un-vignetted      FOV     as $7.8\\arcmin\\times7.8\\arcmin$       (see      OSIRIS      commissioning webpage\\footnote{http://www.gtc.iac.es/en/pa\\-ges\\-/\\-ins\\-tru\\-men\\-ta\\-tion\\-/\\-o\\-si\\-ris\\-/\\-da\\-ta\\--\\-co\\-mmi\\-ssio\\-ning.php}). OSIRIS specifications are extensively described in the manual provided by      the       instrument      team      (OSIRIS       TF      user manual\\footnote{http://www.gtc.iac.es/en/pages/\\-ins\\-tru\\-men\\-ta\\-tion\\-/\\-o\\-si\\-ris\\-/\\-o\\-si\\-ris2\\-.php\\-$\\#$U\\-se\\-ful\\-$\\_$Do\\-cuments; current version dated June 28, 2009} -- this document is referred to along the paper  as OSIRIS {\\em manual}).  In  particular, it provides the wavelength calibration procedure for radii outside the MF.  The OSIRIS FPI is an ideal  instrument to map emission line regions in galaxies.   It provides  a  high throughput  with narrow  transmission bandpass.   The central wavelength  is adjustable  to account  for the galaxy recession velocities, and the field of view (FOV) is reasonably large  with  a  small  pixel  scale  (0\\farcs125).   These  properties together with  the collecting power of  a 10 m  telescope, designed to get the best possible natural image quality, prompted us to design the Local                          Universe                         Survey (LUS\\footnote{http://www.inaoep.mx/$~$gtc$-$lus/}), a research program optimized for OSIRIS.  LUS aims at studying the star-formation history of a  complete sample of nearby  galaxies.  We planned  to obtain high signal-to-noise,  high angular  resolution,  narrow band  maps of  all galaxies  inside a  volume  of  3.5 Mpc  radius.   All irregulars  and spirals inside a volume of 11  Mpc were also included, together with a sample  of  Virgo cluster  galaxies.   These  observations suffice  to produce  a detailed  unique  description of  the  gaseous and  stellar content  of  nearby  galaxies.    LUS  has  been  defined  within  the ESTALLIDOS\\footnote{http://estallidos.iac.es/estallidos/Estallidos.jsp} project  in  close  collaboration   with  the  Instituto  Nacional  de Astrof\\'isica,  \\'Optica  y   Electr\\'onica  in  Mexico.   Significant preparatory work was carried out  in order to prove the feasibility of LUS;  using the  OSIRIS  features  described in  the  manual, we  made simulations to  estimate the exposure  times, to choose  precisely the filters (broad band and TF),  and to optimize the observing procedure. Before submitting LUS as an  ESO-GTC large program, we decided to test our simulations and procedures with real data.   The object selected  to calibrate  the photometric  accuracy was M101, an  extended and well-known nearby galaxy  (7.38~Mpc, Rizzi et al. 2007), that  exceeds the size of OSIRIS  FOV.  Observations were carried out to cover a large  region of the southwestern part of the galaxy,  containing also the  well-known giant  extragalactic \\HII\\, region NGC 5447.   The central wavelengths of the  TF were chosen to map  emission lines  such as  \\Ha\\, or  the sulphur  doublet (\\SII), which  allow  to derive  star-formation  rates  (SFRs) and  electron densities  of  the  star-forming  regions.   Even  if  not  optimal, observations were made using the narrowest bandpass of FWHM 18 \\AA\\, available at the time.  The precise central wavelength of the filter needed   to  separate   \\NII\\  from   \\Ha\\  was   also   tested,  as distinguishing  the  two  lines  is  critical  to  be  able  to  use diagnostic  diagrams  like BPT  \\citep{baldwin81}.   The basic  data reduction  was  carried  out  using  our  own  pipeline.   Then,  we reconstructed monochromatic emission line images taking into account the  wavelength calibration provided  in OSIRIS  documentation. From this analysis,  we found  that \\Ha\\, fluxes  and SFRs  were globally consistent with  values from the literature.   However, the emission line ratios for both \\NII\\  over \\Ha\\ and the sulphur doublet (\\SII) were far from  any observed value and whatever  \\HII\\, region model. In an  attempt to identify the  problem, we started  from scratch by using spectral  calibration lamps  to test wavelength  shifts beyond the MF.  The result of our calibration is the content of this paper. We  have figured  out that  the  prescriptions given  in the  manual suffice for the  MF, while the new calibration  in this paper allows for   using   OSIRIS   in   its   full   capability.    It   permits spectrophotometry in the full FOV with a wavelength accuracy as good as that  in the  MF.  In  our view OSIRIS  is an  extremely powerful instrument which  could be seriously limited  if only the  MF can be used  for  spectrophotometry.  Although  our  work  was designed  to recalibrate our  fields and being able to  recover reliable emission line fluxes over the entire OSIRIS FOV, it is of interest beyond the original scope, and may broaden the community of OSIRIS users.  The  paper  is  organized  as follows:  in  \\S~\\ref{sec:optics}  we introduce the  basic optical concepts required to  understand both the nature    of   the    problem   and    the    calibration   procedure. \\S~\\ref{sec:urge}  shows how  the  original wavelength  calibration needs refinement to provide an accuracy better than 1\\AA\\, in the full FOV.  Based  on the data  introduced in \\S~\\ref{sec:data},  the new calibration    is    worked    out   in    \\S~\\ref{sec:cal}.     In \\S~\\ref{sec:protocol}   we  recommend   a   wavelength  calibration protocol that goes  all the way from a detailed  test of the long-term stability of the system, to a black-box recipe for urgent calibration. The     conclusions    of     the    work     are     summarized    in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  OSIRIS TF and GTC represent a unique combination to study the physical properties  of extended  astronomical  sources.  They  are capable  of providing high S/N, narrow band  images, over a relatively large field with  an excellent  spatial sampling.   The OSIRIS  TF is  based  on a conventional  FPI mounted  in collimated  beam. FPI  based instruments provide  an area  of the  FOV,  the MF,  where the  wavelength can  be considered constant within the  band-pass of the TF.  However, outside the MF there is a significant shift in wavelength which depends on the distance to  the optical  axis.  Therefore, in  order to  benefit from OSIRIS large FOV, an accurate wavelength calibration is needed.  Motivated by inconsistencies in the measurements of line ratios during an observational campaign, we  decided to re-calibrate empirically the equation which describes the shift  in wavelength with position on the FOV.   Using spectral calibration  lamps, we  found that  the proposed nominal calibration suffices only in  the central 2\\arcmin of the FOV, but  fails elsewhere.  The  approximate equation  originally given  to describe the  wavelength shift  with radius is  not correct  for large radii,    and   the    exact   equation    must   be    used   instead (Eq.~(\\ref{eq:caleq1})).  We  identify  three  different  sources  of error  in  the  wavelength calibration,    namely,    errors    in   the    central    wavelength $\\delta\\lambda(0)$, errors  in the calibration  constant $\\delta C^*$, and  errors in  the center  position  $\\delta x_{\\rm  c}$ and  $\\delta y_{\\rm c}$.  The  error in the central wavelength  is intrinsic to the procedure used  to tune the  instrument at the telescope.   We confirm that the  errors introduced  in this process  are $\\lesssim  1$\\AA, as stated  in the  instrument manual.   The calibration  constant  in the procedure proposed here only depends  on the effective focal length of the  camera lens.   It has  been empirically  determined with  a small error (Eq.~(\\ref{eq:focal})),  good enough to  discard the calibration constant as a source of uncertainty.  We have also found a new optical center    defined   independently    for   the    two    OSIRIS   CCDs (Eq.~(\\ref{eq:center})) which  improves significantly the calibration. Actually, the position  of the center seems to  be the factor limiting the  final  accuracy.   Considering  all  uncertainties,  the  updated calibration provides an accuracy of $\\Delta\\lambda \\lesssim 1$\\AA\\, in the entire OSIRIS FOV.  The  use  of  the  new  calibration makes  OSIRIS  ideal  for  studies involving  spectrophotometric measurement  of large  galaxies covering the  full FOV,  such as  those that  triggered the  present  work (see \\S~\\ref{sec:introduction}).   However,  it  is  of  broader  interest. Whatever  study  using  the  entire  FOV will  benefit  from  the  new calibration (e.g.,  extended planetary  nebulae, clusters of  stars or galaxies distributed in a large area, etc).  Based in our previous experience, a wavelength calibration protocol is proposed  in  \\S~\\ref{sec:protocol}.   It  will  allow  to  check  the stability of both OSIRIS center and the calibration constant, i.e., to see  whether  they vary  over  long  periods,  when OSIRIS  optics  is re-aligned,  or  if different  wavelength  ranges  are  used.  In particular, when the blue etalon will be at work, a similar study on the wavelength calibration must be done."
    },
    "1107/1107.0529_arXiv.txt": {
        "abstract": "NGC~4013 is a nearby Sb edge-on galaxy known for its ``prodigious'' \\hi\\ warp and its ``giant'' tidal stream. Previous work on this unusual object shows that it cannot be fitted satisfactorily by a canonical thin+thick disk structure. We have produced a new decomposition of NGC~4013, considering three stellar flattened components (thin+thick disk plus an extra and more extended component) and one gaseous disk. All four components are considered to be gravitationally coupled and isothermal. To do so, we have used the $3.6\\mu{\\rm m}$ images from the Spitzer Survey of Stellar Structure in Galaxies (S$^4$G).  We find evidence for NGC~4013 indeed having a thin and a thick disk and an extra flattened component. This smooth and extended component (scaleheight $z_{\\rm EC}\\sim3$\\,kpc) could be interpreted as a thick disk or as a squashed ellipsoidal halo and contains $\\sim20\\%$ of the total mass of all three stellar components. We argue it is unlikely to be related to the ongoing merger or due to the off-plane stars from a warp in the other two disk components. Instead, we favor a scenario in which the thick disk and the extended component were formed in a two-stage process, in which an initially thick disk has been dynamically heated by a merger soon enough in the galaxy history to have a new thick disk formed within it. ",
        "introduction": "Thick disks, first detected by Burstein (1979) and Tsikoudi (1979), are seen in edge-on galaxies as excesses of light a few thin disk scaleheights above the galaxy mid-planes. They are known to be ubiquitous (Yoachim \\& Dalcanton 2006; Comer\\'on et al.~2011a) and their properties give important clues for understanding galaxy formation and evolution (Comer\\'on et al.~2011b and references therein; hereafter CO11b). Recent studies (Robertson et al.~2006; Elmegreen \\& Elmegreen 2006; Brook et al.~2007; Richards et al.~2010) suggest an in situ formation mechanism for a significant fraction of the thick disk mass. This was recently supported further by our result that the masses of thick and thin disks are of the same order (CO11b). Other formation mechanisms such as disk kinematical heating due to its own overdensities (Villumsen 1985; H\\\"anninen \\& Flynn 2002; Sch\\\"onrich \\& Binney 2009; Bournaud et al.~2009) and the accretion of stars from infalling satellites (Statler 1988; Gilmore et al.~2002; Abadi et al.~2003; Navarro et al.~2004; Martin et al.~2004; Read et al.~2008) also contribute to the thick disk mass (CO11b).  In CO11b we made thin+thick disk decompositions of 46 edge-on galaxies using images from the Spitzer Survey of Stellar Structure in Galaxies (S$^4$G; Sheth et al 2010). However, two galaxies, ESO~079-003 and NGC~4013, could not be successfully fitted down to low surface brightness due to the presence of an extra light component not accounted for in our fits [affecting significantly the luminosity profiles starting at $\\mu_{3.6\\mu{\\rm m}}({\\rm AB})=23\\,{\\rm mag\\,arcsec^{-2}}$]. In CO11b, NGC~3628 also presents an obvious third component, but at a lower surface brightness [$\\mu_{3.6\\mu{\\rm m}}({\\rm AB})=24.5\\,{\\rm mag\\,arcsec^{-2}}$]. The aim of this Letter is to study the properties of these extra components. As ESO~079-003 is located near a bright star which makes its study at low surface brightness levels difficult and NGC~3628 presents a disturbed morphology due to a recent interaction we focused on NGC~4013.  NGC~4013 is an Sb galaxy (Buta et al.~2007), at $18.6\\pm2.5$\\,Mpc (NED average for 13 redshift-independent measurements). Its optical diameter is $D_{25}=294\\arcsec$ (HyperLEDA; Paturel et al.~2003). It is often described using superlative adjectives: it has a ``prodigious'' \\hi\\ warp (Bottema et al.~1987; Bottema 1995; 1996) and a ``giant'' stellar tidal stream with an age of a few Gyr (Mart\\'inez-Delgado et al.~2009; hereafter MD09). The \\hi\\ warp starts just at the optical edge of the galaxy (Bottema~1995) and it is one of the largest warps ever observed. It may have been triggered by the minor merger event which caused the tidal stream (MD09). In addition, NGC~4013 has a boxy bulge (Jarvis 1986) which could be tracing an edge-on bar (Athanassoula \\& Misiriotis 2002; Mart\\'inez-Valpuesta et al.~2006) or be the consequence of a merging event (Binney \\& Petrou 1985).  The S$^4$G $3.6\\,\\mu{\\rm m}$-band image used for the study presented in this Letter is shown in Fig.~\\ref{NGC4013}.  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=0.45\\textwidth]{NGC4013.eps}\\\\ \\includegraphics[width=0.45\\textwidth]{NGC4013_2.eps}\\\\ \\end{tabular} \\caption{\\label{NGC4013} $3.6\\mu{\\rm m}$-band S$^4$G image of NGC~4013 with two different luminosity stretches. The vertical lines indicate the limits of the bins for which luminosity profiles have been produced, from left to right, at galactocentric distances of $-0.8\\,r_{25}<R<-0.5\\,r_{25}$, $-0.5\\,r_{25}<R<-0.2\\,r_{25}$, $-0.2\\,r_{25}<R<0.2\\,r_{25}$, $0.2\\,r_{25}<R<0.5\\,r_{25}$ and $0.5\\,r_{25}<R<0.8\\,r_{25}$. In order to avoid the influence of the bulge we have not produced fits for the central bin.} \\end{center} \\end{figure} ",
        "conclusions": "NGC~4013 is not adequately described by the canonical thin+thick disk description (CO11b). In this Letter, the luminosity profiles of NGC~4013 are fitted satisfactorily using the solutions of three stellar flattened components plus one gaseous disk in equilibrium. The newly described extended component (EC) has a relatively low surface brightness, but due to its vertical extent, contains a significant fraction of the disk mass (between 20\\% and 26\\% depending on the assumed $\\Upsilon_{\\rm t}/\\Upsilon_{\\rm T}$). The EC has a longer scalelength than the galaxy disks, is smooth and its properties do not depend strongly on varying galactocentric radii.  The nature of the EC is unknown and could be a second thick disk or a squashed elliptical halo. The smoothness of the EC makes it unlikely to be related to the ongoing minor merger of NGC~4013. We also discard the EC to be made of off-plane stars of a warped disk, since the warp has been modeled to have its line of nodes in the direction of the line of sight (Bottema 1996). We favor a scenario in which the EC was formed in a two-stage process, in which an initially thick disk was dynamically heated by a merger soon enough in the galaxy history to have a new thick disk formed within it."
    },
    "1107/1107.0974_arXiv.txt": {
        "abstract": "We investigate the origin of the colour-magnitude relation (CMR) observed in cluster galaxies by using a combination of a cosmological {\\em N}-body simulation of a cluster of galaxies and a semi-analytic model of galaxy formation.  The departure of galaxies in the bright end of the CMR with respect to the trend denoted by less luminous galaxies could be explained by the influence of minor mergers. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.0659_arXiv.txt": {
        "abstract": "{}{}{}{}{} \\abstract {We present the analysis of \\xmm and \\swift optical-UV and X-ray observations of the Seyfert-1/QSO \\object{Mrk 509}, part of an unprecedented multi-wavelength campaign, investigating the nuclear environment of this AGN. The \\xmm data are from a series of 10 observations of about 60 ks each, spaced from each other by about 4 days, taken in Oct-Nov 2009. During our campaign, Mrk 509 was also observed with \\swift for a period of about 100 days, monitoring the behaviour of the source before and after the \\xmm observations. With these data we have established the continuum spectrum in the optical-UV and X-ray bands and investigated its variability on the timescale of our campaign with a resolution time of a few days. In order to measure and model the continuum as far as possible into the UV, we also made use of Hubble Space Telescope (\\hst) Cosmic Origin Spectrograph (COS) observations of Mrk 509 (part of our coordinated campaign) and of an archival Far Ultraviolet Spectroscopic Explorer (\\fuse) observation. We have found that in addition to an X-ray power-law, the spectrum displays soft X-ray excess emission below 2 keV, which interestingly varies in association with the thermal optical-UV emission from the accretion disc. The change in the X-ray power-law component flux (albeit smaller than that of the soft excess), on the other hand, is uncorrelated to the flux variability of the soft X-ray excess and the disc component on the probed timescale. The results of our simultaneous broad-band spectral and timing analysis suggest that, on a resolution time of a few days, the soft X-ray excess of Mrk 509 is produced by the Comptonisation of the thermal optical-UV photons from the accretion disc by a warm (0.2 keV) optically thick (${\\tau \\sim 17}$) corona surrounding the inner regions of the disc. This makes Mrk 509, with a black hole mass of about 1--3 $\\times 10^8 \\ \\mathrm{M}_{\\odot}$, the highest mass known system to display such behaviour and origin for the soft X-ray excess.} ",
        "introduction": "Variability studies of Active Galactic Nuclei (AGN) over the last few decades have been essential in developing our understanding of the physical structure of emitting regions in the vicinity of the central engine. The huge amount of energy radiated away is explained by the current paradigm of accretion of material onto a supermassive black hole (SMBH). The emitting regions in the vicinity of the SMBH are too small to be resolved with current telescopes and therefore clues to the processes going on in the central regions come from analysis of their flux and spectral variability. Multi-wavelength monitoring of the AGN variability is a most effective way to investigate complex phenomena such as the accretion onto a SMBH and provides diagnostics of the processes in the nuclear environment of AGN.  The most prominent feature of AGN Spectral Energy Distributions (SEDs) is the so-called `big-blue-bump' peaking in the extreme UV, which was first attributed to thermal emission from the accretion disc by \\citet{Shi78}; it is thought that higher energy photons (UV) originate from small radii in the disc where the temperature is higher, while optical photons are produced further out. The first X-ray spectrum of a Seyfert-1 galaxy, which established the power-law shape of the X-ray continuum, was obtained for NGC 4151, using the proportional counter onboard the {\\it Ariel-V} satellite, by \\citet{Ive76}. The power-law component in the X-ray spectra of AGN is generally believed to be the result of Compton up-scattering of optical-UV photons from the disc by energetic electrons in a hot corona surrounding it. Apart from the power-law, the X-ray continuum often comprises another component, the soft X-ray excess, which was first discovered in the spectrum of Mrk 509 by \\citet{Sin85} using data from the first mission of High Energy Astronomy Observatories (\\heao). Shortly afterwards, the presence of a soft excess was also reported in another Seyfert-1 galaxy, \\object{Mrk 841}, by \\citet{Arn85} using European X-ray Observatory Satellite (\\exosat) observations. Since then, the soft excess has been observed in the X-ray spectra of the majority of AGN and its origin remains an area of active research to this date. Different interpretations for the nature of the soft X-ray excess can be found in the literature such as: (1) the high energy tail of the thermal emission from the accretion disc (e.g. \\citealt{Arn85}, \\citealt{Pou86}); (2) an artefact of strong, relativistically smeared, partially ionized absorption in a wind from the inner disc (e.g. \\citealt{Gie04}); (3) part of the relativistically blurred photoionised disc reflection spectrum modelled by \\citet{Ros05} (e.g. \\citealt{Cru06}); (4) `warm' Comptonisation: up-scattering of seed disc photons in a Comptonising medium which has lower temperature and higher optical depth than the corona responsible for the X-ray power-law emission (e.g. the study by \\citealt{Mag98} of the Seyfert-1 \\object{NGC 5548}; the analyses by \\citealt{Mid09} and \\citealt{Jin09} of two Narrow Line Seyfert-1s (NLS1s): RE J1034+396 and RX J0136.9-3510).  Past multi-wavelength monitoring campaigns of Seyfert-1 galaxies have shown striking similarities in the optical-UV and X-ray variability, suggesting a strong link between the emission in these energy ranges, such as that provided by inverse Comptonisation. \\citet{Wal93} studied the soft X-ray (0.1--2.4 keV) spectra of 58 Seyfert-1 AGN, including Mrk 509, using \\rosat observations. They reported presence of a soft X-ray excess above the extrapolation of the hard X-ray power-law in 90\\% of the sources. They also found that the soft X-ray excess is well correlated to the strength of the big-blue-bump observed with the International Ultraviolet Explorer (\\iue) in the UV band. They concluded that the big-blue-bump is an ultraviolet to soft X-ray feature, which has a similar shape in all Seyfert-1 galaxies of their sample. From an intensive multi-instrument campaign of \\object{NGC 4151}, \\citet{Ede96} found that optical-UV and 1--2 keV X-ray fluxes varied together on timescale of days, with the X-ray flux varying with much larger amplitude than the optical-UV; the phase difference between UV and 1--2 keV X-ray was consistent with zero lag, with an upper limit of $\\lesssim 0.3$ days. \\citet{Mar97} measured the spectrum and lightcurve of \\object{NGC 5548} over a period of 2 months using the Extreme Ultraviolet Explorer (\\euve) when the galaxy was also monitored with \\hst. They found the optical-UV and EUV variations to be simultaneous with the amplitude in the EUV twice that in the UV. Also the shape of the \\euve spectrum was consistent with a gradual decreasing of flux from the UV through to the soft X-ray with no emission lines detected in the EUV band. Furthermore, from a monitoring study of Mrk 509 using Rossi X-ray Timing Explorer (RXTE) and ground-based optical observations, \\citet{Mar08} report a strong correlation between variability of the optical and hard X-ray emission on timescale of a few years, with optical variations leading the hard X-rays by about 15 days.  The work presented here is part of a large multi-wavelength campaign studying the nearby Seyfert-1/QSO Mrk 509. This AGN has a cosmological redshift of 0.034397 (\\citealt{Huc93}) corresponding to a luminosity distance of 145 Mpc (taking ${H_{0}=73\\ \\mathrm{km\\ s^{-1}\\ Mpc^{-1}}}$, $\\Omega_{\\Lambda}=0.73$ and $\\Omega_{m}=0.27$). The details of our multi-wavelength campaign are presented in \\citet{Kaa11}. As part of this campaign, \\swift monitoring was carried out before and after a series of observations by \\xmm and the International Gamma-Ray Astrophysics Laboratory (\\integral). One of the aims of this campaign is determining the location and structure of the multi-phase warm absorber outflows in Mrk 509, which is done by monitoring the response of the warm absorber to intrinsic source variability using high resolution spectra from the Reflection Grating Spectrometer (RGS) onboard  \\xmm. In this work we focus particularly on the optical-UV and X-ray source variability of Mrk 509 using data from the \\xmm Optical Monitor (OM) and the European Photon Imaging Camera (EPIC-pn). Our goal in this paper is to try and explain the intrinsic multi-wavelength variability of Mrk 509 and its relation to the soft X-ray excess in terms of physical changes in the accretion disc and corona.  The structure of this paper is as follows. In Section 2 we describe the observations and data reduction, Sect. 3 focuses on the photometric and spectroscopic corrections applied to the data in order to establish the continuum; the spectral modelling is described in detail in Sects. 4 and 5; we discuss our findings in Sect. 6 and give concluding remarks in Sect. 7. ",
        "conclusions": "\\label{conclusions}  In order to study the intrinsic optical-UV and X-ray continuum of the Seyfert-1/QSO hybrid Mrk 509 (using \\xmm OM and EPIC-pn, \\swift UVOT and XRT, \\hst/COS and \\fuse), we have taken into account various processes occurring in our line of sight towards the central source, such as Galactic reddening and absorption, host galaxy emission, BLR and NLR emission and the AGN multi-phase warm absorption. We conclude that:  \\begin{enumerate}  \\item The optical-UV continuum of Mrk 509 is consistent with thermal emission from a geometrically thin, optically thick cold accretion disc with a maximum disc temperature of about 2 eV. The X-ray continuum can be represented by two distinct components: a power-law with $\\Gamma \\sim 1.9$ and a soft X-ray excess below 2 keV.  \\item We have investigated the variability of the optical-UV and X-ray continuum over the 100 days duration of our monitoring campaign with a resolution of a few days. We have found that the variability of the soft X-ray excess is very similar in profile to that of the optical-UV emission from the disc. The flux variability of the higher energy X-ray power-law component (albeit smaller), on the other hand, is uncorrelated with the variability of the soft X-ray excess and that of the disc emission over the probed timescales. There is however a correlation between the power-law photon index and the disc blackbody flux.  \\item We can account for the origin of the soft X-ray excess and its variability by Comptonisation of the disc photons in a warm (about 0.2 keV), optically thick ($\\tau \\sim 17$) corona surrounding the inner disc. The X-ray power-law is likely to originate from Comptonisation in a separate optically-thin hot corona, possibly located further out, in the form of an atmosphere surrounding the disc.  \\item Explaining the soft X-ray excess using warm Comptonisation has previously been suggested for the Seyfert-1 {NGC 5548} and the NLS1s: {RE J1034+396} and {RX J0136.9-3510}. This interpretation is analogous to the inner-disc warm-corona model which explains the existence of the high-energy tail observed in some BHBs spectra in the Very High State. Mrk 509 is the highest mass black hole system known to display such variability and origin for the soft excess. Only variability studies based on multi-wavelength monitoring campaigns, such as the one reported in this paper, allow the detailed investigations that can reveal the processes leading to the formation of AGN broad-band spectra.  \\end{enumerate}"
    },
    "1107/1107.1696_arXiv.txt": {
        "abstract": "A model is presented to estimate the fraction of Supernova Type-II events (SNII) occurring inside vs. outside a spiral arm for a given star formation episode.  The probability distribution function (PDF) for this fraction is given for use in models similar to those of Shaviv et al. \\cite{Shaviv_2009}\\cite{Piran_2009}. The calculated PDF for the SNII fraction, $SNII_{in/total}$, defined as the number of SNII inside a spiral arm divided by the total number of SNII from a star formation event, provides a constraint on the magnitude of supernova remnant (SNR) concentrations used in cosmic ray propagation models attempting to explain the PAMELA anomaly. Despite the concentration of star formation within spiral arms, this model predicts the majority of SNII events actually occur in inter-arm regions and calls into question the SNR concentration assumption of Shaviv et al. ",
        "introduction": "Recently, the PAMELA collaboration released their results for the Galactic cosmic ray (GCR) positron fraction, $J_{e^{+}}/(J_{e^{-}}+J_{e^{+}})$\\cite{PAMELA_2009}.  The remarkable feature in the PAMELA results is an increase in the positron fraction at energies above $\\sim10$ Gev. Standard GCR propagation models assume all positrons are secondarily produced via interactions with the interstellar medium (ISM). This results in a decreasing positron fraction with energy.  Many models have been proposed to account for this discrepancy \\cite{PAMELA2_2009}: DM particle annihilations, reacceleration of secondaries in supernova remnants (SNRs) (eg. \\cite{Ahlers_2009}), inhomogeneities in the distribution of SNRs (eg. \\cite{Shaviv_2009}\\cite{Piran_2009}). For a more thorough discussion of proposed models and their implications, see \\cite{Fan_2010}. Models relying on a concentration of SNR in the spiral arms of the Milky Way (MW) relative to the inter-arm regions have reproduced many observed results \\cite{Piran_2009}\\cite{Shaviv_2009}. However, the literature is sparse on the specific value of this concentration; thus, it has been difficult to constrain the magnitude of the concentrations used to plausible ranges.  In section 3, I present the model used to calculate $SNII_{in/total}$ as well as the necessary input parameters. Section 4 presents the results of the calculations and the PDF of $SNII_{in/total}$. Section 5 provides discussion on the use of the results and possible refinements. ",
        "conclusions": "Models utilizing a concentration of SNII in the spiral arms to account for the PAMELA anomaly can be roughly tested by the integration of Eqn. 1 for the range of concentration values which allow the model to fit observed data. This provides a simple constraint on the likelihood of the concentration required by the model.  Piran et al. \\cite{Piran_2009}  assume a value of 4 times more SNII in spiral arms than out (4:1). This corresponds to a $SNII_{in/total}$ fraction as used in this paper of 0.8.  The robustness of their model was not discussed with respect to this assumption, but if we assume a range of 3:1 to 5:1 still reproduces observations within their model then we can assess the likelihood.  A concentration between 3:1 and 5:1 ($\\frac{3}{4} \\le SN_{in/total} \\le \\frac{5}{6}$) corresponds to an area under the PDF of $\\sim0.029$. This shows that a ratio of 4 times more SNII in a spiral arm than out is, in this model, not a robust assumption.  Further, this illustrates a general prediction of this SNII model. Namely, for a significant fraction of the parameter space, there will, in fact, be a deficiency of SNII in spiral arms; most SNII will occur in inter-arm regions.  The median calculated value is $\\sim0.329$, corresponding to roughly 2 SNII in the inter-arm region for every 1 SNII inside a spiral arm.  The inputs used in this model are well constrained observationally, with the exception of spiral arm size. The physical extent of spiral arms is inherently variable; further, the edges are not sharply delineated. Coupled with possible eccentricities in the orbits of stars and molecular clouds, the range of values for the effective size of the spiral arm used in the calculations represents the only significant source of systematic error. Clearly, in the model smaller arm sizes directly bias results toward lower $SNII_{in/total}$ fractions with larger arm sizes tending in the opposite direction.  In this paper, the large range in spiral arm size (see Table 1) was used to account for these issues. However, with limits on the overdensity in the ISM required by the models of Piran\\cite{Piran_2009} and Shaviv\\cite{Shaviv_2009}, a better definition of spiral arm size as applied to calculating $SNII_{in/total}$ could be constructed. At that point, variances in orbital eccentricities at the Solar Galactocentric Radius, though relatively small, could be more readily accounted for to construct a more locally applicable version of this general model.  With the majority of SNII occurring outside of the progentior star's parent spiral arm, models attempting to fit the PAMELA anomaly with spatial concentrations of SNII (rather than known locations of individual SNII remnants) must take note of the level of concentration needed and it's  likelihood."
    },
    "1107/1107.3370_arXiv.txt": {
        "abstract": "We report on a detailed investigation of the $\\gamma$-ray emission from 18 broad line radio galaxies (BLRGs) based on two years of {\\it Fermi} Large Area Telescope (LAT) data. We confirm the previously reported detections of 3C~120 and 3C~111 in the GeV photon energy range; a detailed look at the temporal characteristics of the observed $\\gamma$-ray emission reveals in addition possible flux variability in both sources. No statistically significant $\\gamma$-ray detection of the other BLRGs was however found in the considered dataset. Though the sample size studied is small, what appears to differentiate 3C~111 and 3C~120 from the BLRGs not yet detected in $\\gamma$-rays is the particularly strong nuclear radio flux. This finding, together with the indications of the $\\gamma$-ray flux variability and a number of other arguments presented, indicate that the GeV emission of BLRGs is most likely dominated by the beamed radiation of relativistic jets observed at intermediate viewing angles. In this paper we also analyzed a comparison sample of high accretion-rate Seyfert 1 galaxies, which can be considered radio-quiet counterparts of BLRGs, and found none were detected in $\\gamma$-rays. A simple phenomenological hybrid model applied for the broad-band emission of the discussed radio-loud and radio-quiet type 1 active galaxies suggests that the relative contribution of the nuclear jets to the accreting matter is $\\geq 1\\%$ on average for BLRGs, whilst $\\leq 0.1\\%$ for Seyfert 1 galaxies. ",
        "introduction": "A long debated problem in our understanding of accreting supermassive black holes (SMBHs) in the Universe is the unification of different types of active galactic nuclei (AGN) in a framework ascribing their observed diversity to a relatively few differing parameters and factors. For example, it has been widely argued that the difference between type 1 and type 2 AGN --- i.e., those possessing and lacking broad permitted emission lines in their nuclear spectra, respectively --- may be explained by simple geometrical effects involving anisotropic obscuration of the active center viewed at different inclination with respect to the accretion disk axis \\citep{ant93}. The accretion rate was claimed to play a role in this context as well, since `standard' geometrically-thin optically-thick accretion disks formed at high accretion rates \\citep[$> 1\\%$ Eddington;][]{sha73} on the one hand, and radiatively-inefficient geometrically-thick accretion flows expected to be present at low accretion rates \\citep[$< 1\\%$ Eddington;][]{nar94,nar95,abr95} on the other hand, can account for different emission spectra and power outputs of high- and low-power AGN (see an early discussion on this issue by \\citealt{fab95} and \\citealt{las96}, and the more recent one in \\citealt{ghi09}). The other relevant parameter may be also the mass of the SMBH itself, which seems to determine some physical differences between Narrow-Line Seyfert 1 galaxies and `regular' Seyfert 1s \\citep[e.g.,][]{pog00,kom08}.  None of the aforementioned factors can, however, account for the dichotomy between radio-loud AGN (which have jets) and radio-quiet AGN (which don't). One may therefore seek the fundamental difference between AGN producing luminous radio-emitting outflows and those lacking such outflows in yet another parameter of the accretion disk/black hole system. One possibility is that, for example, radio-loud AGN harbor rapidly spinning SMBHs with rotational energy extracted electromagnetically via the \\citet{bla77} mechanism, and converted to the kinetic luminosity of relativistic jets. At the same time, the negligible angular momentum of SMBHs hosted by radio-quiet AGN precludes the formation of such powerful well-collimated outflows. This so-called `spin paradigm' \\citep{bla90} has recently been claimed to be supported, after some modifications, by several observational findings and theoretical investigations \\citep[e.g.,][see also in this context Garofalo 2009]{koi02,sik07,tch10}. Regardless of this debate, the presence of a relativistic jet constitutes a fundamental distinction between various types of AGN, simply because an anisotropic and strongly Doppler-boosted jet emission can dramatically affect the observed properties of a source. In this context, geometrical effects play again a major role. In particular, the total radiative output of those radio-loud AGN which are observed at small viewing angles to the jet axis (`blazars') may be dominated by the broad-band jet emission, while those AGN which are inclined at larger viewing angles (e.g., radio galaxies) may display radiative signatures of both outflowing and accreting matter at comparable levels \\citep[see, e.g.,][]{bar89,urr95}. We note that in addition to such geometrical effects, the age of a radio structure (i.e., the time elapsed after the onset of the jet activity in the nucleus) is yet another factor crucial to understanding unification of radio-loud AGN \\citep[see, e.g.,][]{odea98}.   Interestingly, new deep radio surveys indicate that the radio-loudness parameter --- which is defined as a ratio of the jet-related radio flux to the disk-related optical flux, and is considered as a proxy for the jet production efficiency\\footnote{More accurately, the radio-loudness parameter depends not only on the total power of a jet relative to the total power of an accreting matter, but also on their respective radiative efficiencies. Only if these radiative efficiencies are not very different between various types of AGN (and various AGN of the same type), the radio-loudness parameter may be considered as a good proxy for the jet production efficiency.} --- shows a continuous distribution rather than a sharp division between radio-loud and radio-quiet AGN. This applies to AGN hosted by early-type galaxies and accreting at high rates (i.e., quasars), which are typically studied in this context \\citep[e.g.,][]{whi00}, but holds also when other elliptical-hosted AGN (radio galaxies) spanning a wide range of the accretion rate are taken into account \\citep{sik07}. Moreover, unresolved non-thermal radio emission and jet-like structures have been discovered in classes of AGN considered previously as `radio quiet', i.e., Seyfert galaxies, although the jets in such systems are non-relativistic and weak when compared to the jets found in `classical' radio galaxies and quasars \\citep[e.g.,][and references therein]{ulv89,kuk95,the01,mid04}. The distribution of radio-loudness parameter in Seyferts, which are typically hosted by late-type (disk) galaxies is, however, similarly a continuous function of the accretion rate, as demonstrated first by \\citet{ho01} and \\citet{ho02}. Yet, it was pointed out by \\citet{sik07} that there is a substantial difference in the distribution of the jet production efficiency between disk-hosted and elliptical-hosted AGN, with Seyfert galaxies being characterized, at any accretion rate, by the radio-loudness parameters orders of magnitudes smaller than the analogous parameters characterizing elliptical-hosted radio galaxies or quasars. Clearly, more studies regarding the jet-disk connection in different types of AGN are needed to understand the jet launching processes and the physics of active SMBHs in general.   In this context, broad line radio galaxies (BLRGs) seem ideal targets for an in-depth investigation, since this particular class of very radio-loud AGN exhibits both the disk-related (`Seyfert-like') and the jet-related (`blazar-like') radiative signatures in their broad-band spectra. Unlike blazars, the jets in BLRGs are not pointing directly toward the observer, and so the relativistic beaming effects and the related jet dominance are only moderate. Moreover, unlike narrow-line radio galaxies (NRLGs) --- which are believed to be intrinsically similar but simply inclined at systematically larger jet viewing angles --- BLRGs are not generally obscured by large amounts of dust distributed in torus-like structures around the nucleus, and hence radiative properties of the accretion disks and of the circumnuclear gas can be easily accessed in their case. BLRGs show in particular optical/UV continuum and emission-line characteristics very similar to those of luminous Seyfert galaxies, which indicates high accretion rates and the presence of standard Shakura-Sunyaev disks in both classes of objects. Some fundamental differences in the X-ray spectra between BLRGs and high-accretion-rate Seyferts have been however noted. Specifically, even though the observed X-ray/soft $\\gamma$-ray emission of BLRGs seem still dominated by the moderately absorbed emission by the accreting plasma (i.e., disks and disk coronae), rather than by the jets, the $1-100$\\,keV continua of BLRGs are flatter, and their reflection components (as well as the fluorescent Fe K$\\alpha$ lines) weaker than in the case of luminous Seyfert galaxies \\citep[e.g.,][]{mar91,woz98,sam99,sam02,era00,zdz01,bal07}. Because of such differences, several authors in the past speculated about the non-negligible jet contribution to the X-ray emission of BLRGs, diluting the accretion-related radiative output in the X-ray domain \\citep{woz98,era00,gra02}. This idea was subsequently examined in various different approaches using most recent broadband X-ray data obtained with {\\it BeppoSAX} \\citep{gra06,gra07}, {\\it Suzaku} and {\\it Swift} \\citep{kat07,sam09}.   \\citet{gra07} attempted to disentangle the jet and the disk contributions to the X-ray spectra of three of the brightest BLRGs \\citep[the approach first applied to the case of the quasar 3C~273; see][]{gra04}. For simplicity, they assumed that the accretion disks in BLRGs and Seyfert 1 galaxies produce similar emission continua and reprocessing features, and subsequently allowed for a presence of the Doppler-enhanced jet radiation at an arbitrary level. The fits obtained for 3C~120, 3C~390.3 and 3C~382 showed that the data are indeed consistent with a combination of a thermal component (in a first approximation associated with an accretion disk) and a non-thermal component associated with the beamed radiation (due to a jet). \\citeauthor{gra07} concluded however that jets make only minimal contribution to the X-ray continuum emission of BLRGs, in agreement with the previous findings by \\citet{woz98}. More recently, \\citet{sam09} also proposed that BLRGs may be just clustered at one end of the distribution of the X-ray spectral parameters (e.g., photon indices and reflection albedos) characterizing Seyfert galaxies with significant overlap. Interestingly, both \\citet{gra07} and \\citet{sam09} suggested that the emission of the underlying jet may instead dominate the radiative output of BLRGs at high energy $\\gamma$ rays, and in particular in the GeV regime.   With the successful launch of the {\\it Fermi} Gamma-ray Space Telescope, we now have an unprecedented opportunity to study in detail the $\\gamma$-ray emission from different types of extragalactic sources --- not only blazars, but also radio galaxies \\citep{PerA,M87,CenA,MAGN,kat10}, and other classes of AGN as well \\citep[such as Narrow-Line Seyfert 1 galaxies, for example;][]{NLS}. During the first 15-month of the {\\it Fermi} mission, 11 non-blazar AGN have been detected in the GeV photon energy range \\citep{FAGN,MAGN} by the {\\it Fermi} Large Area Telescope \\citep[LAT;][]{atw09,1FGL}. This `misaligned AGN sample' includes seven Faranoff-Riley type I (low-power) radio galaxies, and four Faranoff-Riley type II (high-power, hereafter `FR II') radio sources consisting of two radio galaxies and two steep spectrum radio quasars. The two luminous radio galaxies detected in $\\gamma$ rays are, in fact, X-ray bright objects classified spectroscopically as BLRGs: 3C~120 (FR I radio morphology), detected for the first time with \\F, and 3C~111 (FR II radio morphology), whose $\\gamma$-ray detection was initially reported by EGRET \\citep{har08} and confirmed by \\F. None of the X-ray bright Seyfert 1 galaxies appear to be detected in $\\gamma$ rays thus far.  In this paper we report on a detailed investigation of the $\\gamma$-ray emission from 18 hard X-ray brightest BLRGs, as well as a comparison sample of high accretion-rate Seyfert 1 galaxies selected as their supposed radio-quiet counterparts (in the framework of the AGN unification scheme). Our primary goals are to examine the $\\gamma$-ray properties of BLRGs as potential `$\\gamma$-ray loud' active galactic nuclei, to study the high-energy jet emission in the selected sources, in particular, investigating the relative contributions of the nuclear jet and accretion disk emission in the broad-band spectra of BLRGs. The two years of \\F exposure provides us with a rather deep all-sky survey reaching the flux limit of typically a few $\\times 10^{-12}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ ($95\\%$ confidence level) between the observed photon energies 100\\,MeV and 10\\,GeV. In $\\S$\\,2, we describe the \\F observations and data reduction procedure. The results of the analysis are given in $\\S$\\,3. The discussion and final conclusions are presented in $\\S$\\,4. Throughout this paper, a $\\Lambda$CDM cosmology with $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.73$, and $\\Omega_m = 0.27$ is adopted \\citep{kom09}. ",
        "conclusions": "The analysis of the two year \\F data for the selected 18 X-ray--bright BLRGs confirmed the previously reported detections of the two sources (3C~120 and 3C~111), at the significance levels not much different between the 15-month and 24-month datasets. This may indicate that the observed $\\gamma$-ray emission is variable on months to year timescales, with a relatively low duty cycle, as revealed by a closer inspection of the corresponding lightcurves (Figure\\,1). In fact, the $>100$ MeV $\\gamma$-ray flux of 3C~111 observed with \\F is (3.5$\\pm$1.2)$\\times$10$^{-8}$ ph cm$^{-2}$ s$^{-1}$, which is about $\\sim 20\\times$ smaller than the maximum recorded by EGRET in the same energy range \\citep[64$\\times$10$^{-8}$ ph cm$^{-2}$ s$^{-1}$,][]{har08}, thus suggesting significant variability in the decade timescale between the EGRET and \\F observations. Moreover, we found some hints for $\\gamma$-ray emission from Pictor A (${\\rm TS} = 20$), even though we cannot claim a formal detection at the moment. These results suggest that other BLRGs could be promising candidates for \\F detections in the near future, and that BLRGs in general could potentially constitute a class of $\\gamma$-ray loud AGN. Such a statement is however difficult to quantify, because of the aforementioned variability of the GeV fluxes from 3C~120 and 3C~111 (assuming that both galaxies are representative for the whole class). On the other hand, the observed flux changes from these two detected galaxies imply that, in analogy with blazar sources, the GeV continua of BLRGs are produced predominantly in the inner parts of their nuclear outflows (jets on scales less than tens or hundreds of parsecs).  In fact, as noted in the previous sections, several authors have previously suggested that the GeV radiation from BLRGs, if detected, should most likely be due to non-thermal and beamed blazar-type emission, but only observed at intermediate inclinations \\citep[jet viewing angles, say, $10^{\\circ} \\lesssim \\theta_j \\lesssim 30^{\\circ}$, versus $\\theta_j \\lesssim 10^{\\circ}$ expected for blazars; see][]{gra07,sam09}. Interestingly, unlike in the luminous blazars, the X-ray/soft $\\gamma$-ray spectra of BLRGs (up to the observed photon energies of $\\sim 100$\\,keV) have been argued to be dominated by the thermal radiation of the accreting matter \\citep{woz98,sam99,era00,zdz01}, with only little (if any) jet contribution \\citep{gra07,kat07,sam09}. Hence it is clear that BLRGs are truly ideal targets for investigating the AGN jet-disk connection in the X-ray/$\\gamma$-ray regime. Here we present a few considerations regarding this issue.  \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.47]{Fig7-a.ps} \\includegraphics[angle=0,scale=0.47]{Fig7-b.ps} \\includegraphics[angle=0,scale=0.47]{Fig7-c.ps} \\includegraphics[angle=0,scale=0.47]{Fig7-d.ps} \\caption{The dependence of the $\\gamma$-ray luminosities (or upper limits for such), expressed in the Eddington units ($L_{\\gamma} / L_{\\rm Edd}$) or as a ratio of the GeV and X-ray monochromatic luminosities ($L_{\\gamma} / L_{\\rm X}$), on (i) the black hole mass $\\mathcal{M}_{\\rm BH}$ ({\\it top left} panel), (ii) the accretion ratio $L_{\\rm acc}/L_{\\rm Edd}$ estimated from the SED fitting ({\\it top right} panel), (iii) the proxy of the radio loudness parameter $L_{\\rm R} / L_{\\rm B}$ ({\\it bottom left} panel),  and (iv) the observed proxy of the accretion rate $L_{\\rm B} / L_{\\rm Edd}$ ({\\it bottom right} panel). BLRGs and Seyferts are denoted by blue and red symbols, respectively. Large filled circles represent the two BLRGs detected by \\F (3C~111 and 3C~120), while the medium-size open circles represent Pictor A. In the {\\it bottom left} panel, when total radio fluxes are used instead of the nuclear radio fluxes, BLRGs are denoted by cyan symbols. Similarly, Seyfert 1 galaxies are denoted by magenta symbols when total radio fluxes are used instead of the nuclear radio fluxes. An upper limit is given for Fairall 9. The ratios of the GeV and X-ray monochromatic luminosities ($L_{\\gamma} / L_{\\rm X}$) expected from a hybrid model are indicated as dashed lines in the bottom panel, for different values of the mixing parameter $\\eta$ = 10$^{-2}$, 10$^{-1}$, and 1.} \\end{center} \\end{figure}     Figure\\,3 presents spectral energy distributions (SEDs) of the two aforementioned $\\gamma$-ray--detected BLRGs; the SEDs of all the remaining objects, for which only the upper limits in the \\F range were derived, are given in Figures\\,8 and 9 further below (see Appendix A). In the figures, the \\F data are indicated by red circles (or arrows), black squares represent the historical data from NED, while the magenta squares show the $5$\\,GHz radio fluxes of the unresolved nuclei. Even though the broad-band spectra vary to some degree between 3C~120 and 3C~111 (e.g., with respect to the ratio between the total and the nuclear radio luminosities), one can gauge the main similarities and differences between the BLRG-type and Seyfert-type SEDs. In particular, while the infrared--to--hard X-ray continua of the two BLRGs are very similar to the one characterizing Seyfert galaxies considered here,  the former objects are significantly brighter than the representative Seyfert in the radio and $\\gamma$-ray domains. Such a dramatic difference in the radio loudness is related to the presence of a relativistic jet, as discussed above, and here we suggest that same is true regarding the $\\gamma$-ray loudness. To investigate it further more quantitatively (yet still illustrative), we consider a simple phenomenological `hybrid' model for the broad-band emission of a type 1 AGN, consisting of individual thermal and the non-thermal emission components \\citep[the approach analogous to the one adopted in][]{gra04,gra07}. The thermal component, related to the radiative output of the accreting and circumnuclear matter (accretion disk, disk corona, dusty torus), is approximated here by the template Seyfert spectrum given by \\citet{kor99}. For the non-thermal broad-band emission component of the nuclear (blazar-type) relativistic jet observed at intermediate viewing angles, we use the well-constrained SED of the radio-loud quasar 3C~273 \\citep[from][]{sol08} as a template\\footnote{It should be emphasized here that the observed multiwavelength emission of 3C~273 is \\emph{not purely non-thermal} in origin. For example, \\citet{gra04} estimated that about a quarter to half of the $2-10$\\,keV flux of this source could be due to the accreting matter, and not due to the jet. However, the ``big blue bump'' is clearly visible in the overall SED of the quasar around UV/soft X-ray frequencies is almost certainly due to the accretion disk. Nevertheless, the broad-band spectrum of 3C~273 is the best sampled (and best understood) spectrum of a quasar dominated in radio and in $\\gamma$ rays by the moderately-beamed emission of a powerful jet observed at intermediate viewing angles \\citep[$\\theta_j \\sim 10^{\\circ}$; see][and references therein]{cou98}. As such, it can indeed be consider as the best available template for the jet-related emission of BLRGs.}. Both templates, plotted in Figure\\,3 (as well as in Figures \\,8 and 9) as green and blue curves, respectively, are rescaled to match the data points for the analyzed sources. If the \\F did not detect the source and only upper limits on $\\gamma$-ray flux were derived, its non-thermal component is entirely determined by the 5 GHz nuclear flux (in this context, see Table~3 of \\citet{ghi05}, who predicted $\\gamma$-ray fluxes of 3CR FR I radio galaxies from 5 GHz core fluxes). This allows us to approximate the relative contribution of the non-thermal and thermal emission components to the observed SED of each object, in terms of a `mixing' parameter ($\\eta$). Under the adopted definition, this parameter increases with the increasing contribution of the jet-like component, and equals unity when both the 3C~273 and the Seyfert template spectra provide comparable contributions to the observed UV flux of a source. Note that the direct radiation of the standard (Shakura-Sunyaev) AGN accretion disk is expected to peak in the UV regime (photon energies $\\sim 10$\\,eV).   Obviously, the adopted model is an oversimplification of a realistic situation, as it ignores several complications which may be potentially relevant. For example, it is at some level questionable to use the broad-band spectrum of the radio-loud quasar 3C~273 as a template for the jet-related emission of Seyfert galaxies in general. Moreover, possible (or even likely) temporal variability of each source considered, as well as the {\\it ad hoc} adopted procedure of matching the assumed templates to the collected data points, precludes any robust statistical analysis (e.g., $\\chi^2$ fitting) which would allow for the values and errors of the $\\eta$ parameter to be determined more accurately. Nevertheless, it is worth noting that this simple model seems to match the SEDs of the two BLRGs well from radio to $\\gamma$-ray frequencies, and that in both cases the emerging values of the $\\eta$ parameter are similar ($\\simeq 0.15-0.35$), being in addition consistent with the ones following from the analysis by \\citet{gra07}. Importantly, the model when applied to the other BLRGs analyzed here returns similar values for the $\\eta$ parameter at the level of few tens of percent at most (as determined by the nuclear radio fluxes), without violating the \\F upper limits (see column 8 in Table\\,2 and the SEDs presented in Figures\\,8). Figure\\,4 presents the distribution of the mixing parameter $\\eta$ emerging from the SED matching. This distribution suggests that, even though the particular frequency ranges may be dominated by one of the two main emission components, the \\emph{total} observed luminosities of the nuclear jets in BLRGs constitute on average not less than $1\\%$ of the accretion-related luminosities (median $\\eta \\approx 0.10$), and that in the case of a few particular objects the jet--to--disk luminosity ratio may even approach unity. An interesting implication of the above statement, strengthened thanks to the inclusion of the \\F data in our modeling, is that one should expect a non-negligible non-thermal emission component in BLRGs in the MeV energy range, constituting as much as $\\sim 1\\%-10\\%$ of the total power emitted in the hard X-ray domain. If correct, this may be of importance for understanding the recently debated origin of the extragalactic MeV background \\citep[see the discussions in][]{ino08,aje09}. On the other hand, the contribution of the jet emission to the total radiative output of the Seyfert 1 galaxies  are in general very small (median $\\eta \\approx 5 \\times 10^{-4}$; see Figure\\,4 and the SEDs shown in Figure\\,9).   To investigate further the collected dataset, we plot the optical $[\\nu F_{\\nu}]_{\\rm B}$ versus X-ray $[\\nu F_{\\nu}]_{\\rm 2-10\\,keV}$  fluxes for all the sources included in the sample in Figure\\,5 (the {\\it top} panel), with BLRGs denoted by blue filled circles and Seyferts by red crosses. The {\\it bottom} panel of Figure\\,4 presents instead the radio $[\\nu F_{\\nu}]_{\\rm 5\\,GHz}$ versus X-ray fluxes. Here, BLRGs are plotted as blue filled cycles when total radio fluxes are used, and as cyan open circles when nuclear radio fluxes (obtained predominantly from VLBI observations where available, supplemented by VLA measurements in a few cases) are considered. In both panels, large symbols indicate the two objects detected by \\F (3C~111 and 3C~120), while the medium-size symbols represent Pictor A galaxy which, even though not formally detected, is characterized by the third-highest TS in the analyzed two-year \\F dataset. These flux-flux plots show that the optical and X-ray fluxes of BLRGs are roughly linearly correlated, as expected, though with a substantial (order-of magnitude) scatter. In addition, however, it is revealed that 3C~111, 3C~120, and also Pictor A, are at the same time the brightest in radio, but only when the nuclear (and not the total) radio fluxes are taken into account. Moreover, the nuclear radio fluxes seem well correlated with the X-ray fluxes for the whole BLRG population. Hence one may conclude that \\F detects preferentially those BLRGs which are characterized by the brightest radio and X-ray nuclei. This supports the idea stated previously that the GeV emission of BLRGs is dominated by the innermost parts of their jets, and is therefore `blazar-like,' being dependent on the jet luminosity and viewing angle. As argued above, the quasar 3C~273, for which the GeV luminosity is about 100 times higher than the nuclear radio luminosity, should provide a reasonably accurate template for such an emission. Indeed, as illustrated in Figure\\,6, the $\\gamma$-ray--to--radio energy flux ratios for the two BLRGs detected by \\F are of the order of $\\simeq 100$, while the corresponding upper limits for all the other objects from the sample (including Seyferts) are above or much above this value. This indicates that the detections of BLRGs in $\\gamma$ rays are at present limited by the sensitivity of the \\F instrument. In fact, from visual inspection of Figure 8 and 9, the predicted $\\gamma$-ray flux from the template is very close to the \\F upper limits for Pictor~A, 3C~390.3, 3C~445, B3~0309+411B and PKS 2153-69.  Also it is important to note that the pc-scale components of the radio jets in both 3C~111 and 3C~120 are characterized by apparent superluminal motions with maximum velocities $\\beta_{\\rm app} \\simeq 5.9$ and $5.3$, respectively, from MOJAVE monitoring data \\citep{lis09}. The corresponding upper limits to the jet viewing angles are thus $\\theta_j \\lesssim 20^{\\circ}$. Interestingly, mildly relativistic velocities are inferred for 3C~390.3 and Pictor~A \\citep[$\\beta_{\\rm app}$ up to $\\simeq 2.2$ and $1.6$, respectively;][]{kel04,tin00}, both of which are good candidates for future \\F detections (see Table\\,2 and Figure\\,8). These four BLRGs, not coincidentally, have the brightest radio/VLBI cores in the analyzed sample, consistent also with a relatively large degree of relativistic beaming (see the comment in Appendix\\,A). Note also that 3C~303 and 3C~382 are other BLRGs with possible mildly relativistic pc-scale VLBI jets \\citep{gio01}.  But what else --- besides the nuclear radio jets --- could be the source of high-energy emission in the considered objects? Clearly, large-scale structures (i.e., extended lobes, large-scale jets and terminal hotspots), which are particularly prominent in BLRGs, are expected to contribute \\emph{at least at some level} to the $\\gamma$-ray emission in the GeV band, since the non-thermal synchrotron continua of these structures extend to the highest radio frequencies and even into the X-ray energies in some cases \\citep[see in this context \\F detection of the extended lobes in the low-power radio galaxy Centaurus~A;][]{cena-lobe}. Such structures are on the other hand almost absent in Seyfert galaxies, due to much less prominent jet activity in those objects. Other possible $\\gamma$-ray emission sites in both BLRGs and Seyferts are their accretion disks and disk coronae. Several related (and rather preliminary) studies presented in the literature, even though not directly applicable to the types of astrophysical objects discussed here, indicate that other additional processes than those typically considered for generating high-energy photons within the accreting matter (e.g., proton-proton interactions), may be relevant in this context. In these scenarios, the resulting disk-related $\\gamma$-ray emission may strongly depend on the main parameters of a black hole/accretion disk system \\citep[such as the black hole mass, spin, accretion rate, and disk inclination; see, e.g.,][]{mah97,oka03,nie09}. Hence it is interesting to look into this issue in more detail for all the sources included in our sample.  Figure\\,7 presents therefore the dependence of the $\\gamma$-ray luminosities (or upper limits for such), expressed in the Eddington units ($L_{\\gamma} / L_{\\rm Edd}$) or as a ratio of the GeV and X-ray monochromatic luminosities ($L_{\\gamma} / L_{\\rm X}$), on  (i) the black hole mass $\\mathcal{M}_{\\rm BH}$ ({\\it top left} panel), (ii) the accretion ratio $L_{\\rm acc} / L_{\\rm Edd}$ determined from the model fitting ({\\it top right} panel), (iii) the proxy of the radio loudness parameter $L_{\\rm R} / L_{\\rm B}$ ({\\it bottom left} panel), and (iv) the observed proxy of the accretion rate $L_{\\rm B} / L_{\\rm Edd}$ ({\\it bottom right} panel). In the figure, BLRGs and Seyferts are denoted by blue and red symbols, respectively. Large filled circles represent the two BLRGs detected by \\F (3C~111 and 3C~120), while the medium-size open circles represent Pictor A. In the {\\it bottom left} panel, when nuclear radio fluxes are used instead of the total radio fluxes, BLRGs are denoted by cyan symbols, and Seyferts by magenta symbols. A ratio of the GeV and X-ray monochromatic luminosities ($L_{\\gamma} / L_{\\rm X}$) expected from a hybrid model is given as dashed lines in the {\\it bottom right}  panel, for different values of a mixing parameter $\\eta$ = 10$^{-2}$, 10$^{-1}$, and 1.   Note that the monochromatic luminosity ratio $L_{\\rm R} / L_{\\rm B}$ is simply proportional to the `standard' radio-loudness parameter $\\mathcal{R}$, and in particular that $L_{\\rm R} / L_{\\rm B} = (\\nu_{\\rm R} / \\nu_{\\rm B}) \\times \\mathcal{R} \\sim 10^{-5} \\, \\mathcal{R}$. Also, for the high accretion-rate objects analyzed in this paper one expects the accretion luminosity $L_{\\rm acc} \\simeq L_{\\rm tot} \\simeq 10 \\times L_{\\rm B}$. As shown, the two galaxies which are the brightest in $\\gamma$ rays, and also Pictor~A, are not characterized by any outstanding values of the radio loudness, accretion rate, or black hole mass, but are in fact quite representative for the other BLRGs included in the sample. This supports once again our conclusion that the detected GeV emission of BLRGs is related predominantly to nuclear jets, and not the other possible nuclear emission components. What should be noted for constraining theoretical models regarding the high-energy emission of accretion disks in AGN is the fact that, at least for a number of sources, the \\F upper limits evaluated here already probe the GeV emission of BLRGs and Seyferts down to the levels of $1\\%$ of the X-ray (disk-related) luminosity, or equivalently $0.01\\%$ Eddington luminosity (see Figure\\,7)."
    },
    "1107/1107.4882_arXiv.txt": {
        "abstract": "The existence of a significant flux of antiprotons confined to Earth's magnetosphere has been considered in several theoretical works. These antiparticles are produced in nuclear interactions of energetic cosmic rays with the terrestrial atmosphere and accumulate in the geomagnetic field at altitudes of several hundred kilometers. A contribution from the decay of albedo antineutrons has been hypothesized in analogy to proton production by neutron decay, which constitutes the main source of trapped protons at energies above some tens of MeV. This Letter reports the discovery of an antiproton radiation belt around the Earth. The trapped antiproton energy spectrum in the South Atlantic Anomaly (SAA) region has been measured by the PAMELA experiment for the kinetic energy range 60--750 MeV. A measurement of the atmospheric sub-cutoff antiproton spectrum outside the radiation belts is also reported. PAMELA data show that the magnetospheric antiproton flux in the SAA exceeds the cosmic-ray antiproton flux by three orders of magnitude at the present solar minimum, and exceeds the sub-cutoff antiproton flux outside radiation belts by four orders of magnitude, constituting the most abundant source of antiprotons near the Earth. ",
        "introduction": "The PAMELA collaboration has recently reported the cosmic ray (CR) antiproton spectrum and antiproton-to-proton ratio measurements in the kinetic energy range 60 MeV--180 GeV \\citep{PRL,PRL_BIS}. These data significantly improve those from previous experiments thanks to the high statistical significance and wide energy interval. The results agree with models of purely secondary production where antiprotons are produced through interactions of CRs with the interstellar medium.  Antiprotons are also created in pair production processes in reactions of energetic CRs with Earth's exosphere. Some of the antiparticles produced in the innermost region of the magnetosphere are captured by the geomagnetic field allowing the formation of an antiproton radiation belt around the Earth. The particles accumulate until they are removed due to annihilation or ionization losses. The trapped particles are characterized by a narrow pitch angle\\footnote{The pitch angle is the angle between particle velocity vector and geomagnetic field line.} distribution centered around 90 deg and drift along geomagnetic field lines belonging to the same McIlwain $L$-shell\\footnote{An $L$-shell is the surface formed by azimuthal rotation of a dipole field line. In a dipole, $L$ is the radius where a field line crosses the equator; in case of the Earth dipole, it is measured in units of Earth radii. McIlwain's coordinates \\citep{McIlwain}, $L$ and $B$ (the magnetic field strength), are pairs describing how far away from the equator a point is located along a given magnetic line at the distance $L$ from the Earth.} where they were produced. Due to magnetospheric transport processes, the antiproton population is expected to be distributed over a wide range of radial distances.  According to the so-called CRAND (Cosmic Ray Albedo Neutron Decay) process \\citep{waltcrand,protcrand}, a small fraction of neutrons escapes the atmosphere and decays within the magnetosphere into protons, which become trapped if they are generated with a suitable pitch angle. Such a mechanism is expected to produce antineutrons (through pair production reactions such as $pp\\rightarrow ppn\\bar{n}$) which subsequently decay to produce antiprotons (CRANbarD). This source is expected to provide the main contribution to the energy spectrum of stably trapped antiprotons and the resulting flux is predicted to be up to several orders of magnitude higher than the antiproton flux from direct $p\\bar{p}$ pair production in the exosphere \\citep{Fuki,Selesnick}.  The magnetospheric antiproton flux is expected to exceed significantly the galactic CR antiproton flux at energies below a few GeV. However, predictions differ and suffer from large uncertainties, especially regarding contributions from the CRANbarD process. A measurement by the Maria-2 instrument \\citep{MARIA2} on board the ``Salyut-7'' and ``MIR'' orbital stations allowed an upper limit on the trapped antiproton-to-proton ratio of $5\\cdot 10^{-3}$ to be established below 150 MeV. This Letter describes the first detection of antiprotons trapped in the inner radiation belt, using the PAMELA satellite-borne experiment. ",
        "conclusions": "Antiprotons trapped in Earth's inner radiation belt have been observed for the first time by the PAMELA satellite-borne experiment. The antiparticle population originates from CR interactions in the upper atmosphere and subsequent trapping in the magnetosphere. PAMELA data confirm the existence of a significant antiproton flux in the SAA below $\\sim$ 1 GeV in kinetic energy. The flux exceeds the galactic CR antiproton flux by three orders of magnitude at the current solar minimum, thereby constituting the most abundant antiproton source near the Earth. A measurement of the sub-cutoff antiproton spectrum outside the SAA region is also reported. PAMELA results allow CR transport models to be tested in the terrestrial atmosphere and significantly constrain predictions from trapped antiproton models, reducing uncertainties concerning the antiproton production spectrum in Earth's magnetosphere."
    },
    "1107/1107.3146_arXiv.txt": {
        "abstract": "We present $R$-matrix calculations of photoabsorption and photoionization cross sections across the K edge of the Li-like to Ca-like ions stages of Ni. Level-resolved, Breit-Pauli calculations were performed for the Li-like to Na-like stages. Term-resolved calculations, which include the mass-velocity and Darwin relativistic corrections, were performed for the Mg-like to Ca-like ion stages. This data set is extended up to Fe-like Ni using the distorted wave approximation as implemented by {\\sc autostructure}. The $R$-matrix calculations include the effects of radiative and Auger dampings by means of an optical potential. The damping processes affect the absorption resonances converging to the K thresholds causing them to display symmetric profiles of constant width that smear the otherwise sharp edge at the K-shell photoionization threshold. These data are important for the modeling of features found in photoionized plasmas. ",
        "introduction": "The K lines and edges of iron and nickel are important diagnostics in X-ray astronomy. They are  in a relatively unconfused part of the X-ray spectrum ($\\sim 6$--10 keV), and they are formed over a wide range of physical conditions.  This can range from cold (by X-ray standards) neutral gas at temperatures $\\leq 10^4$ K, to highly ionized gas near $10^8$ K. Their energy and shape give information about the ionization and kinematics in the gas. The strengths of emission features, such as K fluorescence lines, constrain the amount of emitting gas and the elemental abundances, while the strengths of absorption edges constrain the column densities along the line of sight to the background X-ray source. Iron K lines and edges have been studied by many existing and past X-ray astronomy instruments. They have provided insights into diverse topics such as elemental abundances in clusters of galaxies \\citep{serletal77}, the spin of black holes \\citep{brenetal11,mccletal11}, and the geometry of stellar flares \\citep{testetal08}. Nickel features have not yet been used to the same extent, owing to the fact that they are weaker by a factor $\\sim$10 from the typical Ni/Fe abundance ratio, and to the fact that many detectors have collecting areas which decrease rapidly above 7 keV. Nonetheless, Nickel has been detected in systems ranging from black hole X-ray binaries \\citep{milletal06,kalletal09}, clusters of galaxies \\citep{tamuetal09}, the center of our galaxy \\citep{koya11}, and supernova remnants \\citep{tama10}. Further impetus for calculations of the spectra of nickel comes from the {\\em Astro-H} satellite, which will result in the most sensitive observations to date of objects at energies containing K features of nickel.  This is because of both the relatively large sensitivity at the energy of the nickel K lines and edges (7.2 -- 10.7 keV), and because of its spectral resolution, which will exceed that of previous experiments at energies above 3 keV.  Over the last several years, our group has worked on providing accurate K-shell atomic data for all ions of astrophysical interest. We have calculated energy levels, wavelengths, Einstein $A$-coefficients, Auger rates, and photoabsorption/ionization cross sections. Complete data sets have been computed for all ion stages of Fe \\citep{bautetal03,palmetal03,palmetal03b,mendetal04,bautetal04,kalletal04}, O \\citep{garcetal05}, N \\citep{garcetal09}, as well as Ne, Mg, Si, S, Ar, and Ca \\citep{palmetal08,wittetal09,wittetal11}. There has also been recent work on the C sequence by \\citet{hasoetal10} using similar techniques.  The data we present in this work are important both for their direct effect on observed astrophysical X-ray spectra via absorption features, and also through their effect on ionization balance calculations, which in turn affects many of the observables from these elements. In particular, it is the detailed resonance and edge structure, missing in previous calculations, which is most important.  There have been a large number of calculations for these systems going back to the 1920's, but they are confined to the non-resonant background and few span the K edge region. For the ions considered in this work, there is good agreement between our results and the background cross section both above and below the K edge characterized by the fits of \\citet{vernyako95}. Laboratory measurements have been mainly confined to neutral systems for photon energies beyond the K edge. A comprehensive assessment and fitting of measurements has been performed by \\citet{veig73}.  There has been recent progress measuring the resonance structure near the K edge of Li-like systems: C$^{3+}$ \\citep{mulletal09} and B$^{2+}$~\\citep{mulletal10}, as well as Be-like C$^{2+}$ \\citep{sculletal05}. $R$-matrix calculations performed as part of those works are in good agreement with the measurements apart from small discrepancies in some resonance positions and heights. With the advances of free electron lasers to produce X-rays \\citep{kantetal06}, we expect it will be possible over the next several years to perform similar experiments for the K-shell absorption by heavier species. However, such measurements for Ni are not yet available. ",
        "conclusions": "Total photoabsorption and photoionization cross sections have been computed for the Li-like to Fe-like ions of Ni. Level-resolved, Breit-Pauli $R$-matrix calculations were carried out for the Li-like to Na-like ions. Term-resolved $R$-matrix calculations were performed for Mg-like to Ca-like where the mass-velocity and Darwin relativistic corrections are included. Distorted wave calculations for Sc-like to Fe-like Ni were performed using {\\sc autostructure} where resonances are included using the isolated resonance approximation.  For all of the $R$-matrix calculations, total photoabsorption and total/partial photoionization data are available. Radiation and Auger dampings are included to account for radiation losses and the spectator Auger process. No damping is included for the {\\sc autostructure} calculations so we present only total photoabsorption cross sections which are convolved with a Gaussian profile using a width representative of currently deployed X-ray detectors.  We find overall good agreement with the background cross sections given by \\citet{vernyako95}, but our positions of the K edge are typically 40-80 eV higher. For Li-like Ni, we are in good agreement with \\citet{naha05} as to resonance positions. However, we find much stronger resonant enhancement due to the finer energy mesh used in the present work.  All data are provided as on-line tables accompanying this article and can also be obtained through the {\\sc xstar} atomic database \\citep{bautkall01} and the Universal Atomic Database. The data sets provided here together with the energy levels and radiative and Auger rates reported in \\citet{palmetal08b} will help modelers to carry out detailed studies of K spectra of Ni ions."
    },
    "1107/1107.1719_arXiv.txt": {
        "abstract": "{ It is well known that $R$-symmetric models dramatically alleviate the SUSY flavor and CP problems.  We study particular modifications of existing $R$-symmetric models which share the solution to the above problems, \\emph{and} have interesting consequences for electroweak baryogenesis and the Dark Matter (DM) content of the universe.  In particular, we find that it is naturally possible to have a strongly first-order electroweak phase transition while simultaneously relaxing the tension with EDM experiments.  The $R$-symmetry (and its small breaking) implies that the gauginos (and the neutralino LSP) are pseudo-Dirac fermions, which is relevant for both baryogenesis and DM. The singlet superpartner of the $U(1)_{Y}$ pseudo-Dirac gaugino plays a prominent role in making the electroweak phase transition strongly first-order.  The pseudo-Dirac nature of the LSP allows it to behave similarly to a Dirac particle during freeze-out, but like a Majorana particle for annihilation today and in scattering against nuclei, thus being consistent with current constraints.  Assuming a standard cosmology, it is possible to simultaneously have a strongly first-order phase transition conducive to baryogenesis \\emph{and} have the LSP provide the full DM relic abundance, in part of the allowed parameter space.  However, other possibilities for DM also exist, which are discussed.  It is expected that upcoming direct DM searches as well as neutrino signals from DM annihilation in the Sun will be sensitive to this class of models. Interesting collider and Gravity-wave signals are also briefly discussed.}  \\end{center} \\end{titlepage}  \\def\\simgt{\\mathrel{\\lower2.5pt\\vbox{\\lineskip=0pt\\baselineskip=0pt \\hbox{$>$}\\hbox{$\\sim$}}}} \\def\\simlt{\\mathrel{\\lower2.5pt\\vbox{\\lineskip=0pt\\baselineskip=0pt \\hbox{$<$}\\hbox{$\\sim$}}}}  \\renewcommand{\\l}{\\langle} \\renewcommand{\\r}{\\rangle} \\newcommand{\\be}{\\begin{eqnarray}} \\newcommand{\\ee}{\\end{eqnarray}}  \\newcommand{\\dd}[2]{\\frac{\\partial #1}{\\partial #2}} \\newcommand{\\NN}{\\mathcal{N}} \\newcommand{\\LL}{\\mathcal{L}} \\newcommand{\\MM}{\\mathcal{M}} \\newcommand{\\ZZ}{\\mathcal{Z}} \\newcommand{\\WW}{\\mathcal{W}}   \\newcommand{\\FO}{{\\rm FO}} \\newcommand{\\FI}{{\\rm FI}} \\newcommand{\\TFO}{T_{\\rm FO}} \\newcommand{\\TFOp}{T_{\\textrm{FO}'}} \\newcommand{\\YFO}{Y_{\\rm FO}} \\newcommand{\\zFO}{z_{\\rm FO}} \\newcommand{\\FOp}{\\textrm{FO}'} \\newcommand{\\vev}{\\textit{vev}} \\newcommand{\\vevs}{\\textit{vevs}}    \\newcommand{\\sv}{\\langle \\sigma v \\rangle}  \\newcommand{\\MPl}{M_{\\rm Pl}}  \\tableofcontents ",
        "introduction": "Although the Standard-Model is an excellent description of Nature up to energies of around a hundred GeV, it fails to explain two of the most important mysteries of our Universe - the nature of Dark Matter (DM) and the origin of the matter-antimatter asymmetry.~\\footnote{We have nothing to say in this paper about the \\emph{biggest} mystery - the tiny value of the Cosmological Constant.} Many candidates for Dark Matter - axions, gravitinos, WIMPs, etc.  exist; however the most popular among them is the WIMP, which can provide the observed relic abundance from a thermal freeze-out mechanism in a large region of parameter space.  Moreover, such a particle also arises within many schemes of beyond-the-Standard-Model physics such as low-scale supersymmetry, extra dimensions, etc.  Similarly, many popular mechanisms for generating the matter-antimatter asymmetry exist.  These can be broadly divided into those which depend on physics at very high scales and those which depend on physics at the electroweak (EW) scale and below.  Models in the first class include models of GUT baryogenesis~\\cite{Kolb:1983ni} and leptogenesis~\\cite{Fukugita:1986hr}.  On the other hand models in the second class include those of electroweak baryogenesis~\\cite{Kuzmin:1985mm}, TeV-scale leptogenesis~\\cite{Flanz:1996fb} and Affleck-Dine baryogenesis~\\cite{Affleck:1984fy}.  There also exist models which propose to explain both the baryon asymmetry of the universe (BAU) and the DM abundance as an \\emph{asymmetry} in a conserved charge~\\cite{Kaplan:1991ah}-\\cite{Fujii:2002aj}, \\cite{Hooper:2004dc}, although they typically involve a very weakly coupled ``hidden\" sector.  Electroweak baryogenesis is attractive since it only depends on physics in the visible sector at the electroweak scale; hence the mechanism is completely testable (at least in principle) at colliders probing energies of order the electroweak scale, or astrophysical/cosmological observations probing temperatures (or times) of order the electroweak scale.  Supersymmetric extensions of the Standard Model provide an elegant solution to the hierarchy problem as well as a Dark Matter candidate in the form of the lightest supersymmetric particle or LSP (assuming $R$-parity conservation). However, within the minimal supersymmetric Standard Model (MSSM), the allowed region for a strongly first-order phase transition, which is one of the important requirements for electroweak baryogenesis, is severely constrained~\\cite{Carena:1996wj} and will soon be tested by experiments (see \\cite{Carena:2008vj} for a recent review).  In addition, there is tension between having enough CP violation to produce the required baryon asymmetry and constraints from electric dipole-moment (EDM) searches~\\cite{Pospelov:2005pr}.  It is, therefore, worthwhile to study extensions of the MSSM with additional terms in the tree-level scalar potential which could help in making the transition strongly first-order.\\footnote{Alternatively, it has been argued that new fermions, strongly coupled to the SM, can also strengthen the EWPT, and lead to a EWBG/DM connection~\\cite{Carena:2004ha}.  Ideas for generating the baryon asymmetry at the TeV scale \\textit{without} a strongly first-order EWPT have also been proposed~\\cite{Shu:2006mm}.} The viability of singlet extensions of the MSSM with regard to electroweak baryogenesis has been widely studied \\cite{Pietroni:1992in} \\cite{Menon:2004wv}. In particular, Ref.~\\cite{Menon:2004wv} studied a singlet extension of the MSSM called the nMSSM, in which it was shown that it was possible to have a strongly first-order phase transition as well as generate the correct relic abundance for the lightest neutralino.  The lightest neutralino in that case is, however, not completely stable but has a lifetime longer than the age of the Universe.  In this work, we explore a different extension of the MSSM which can give rise to electroweak baryogenesis and a WIMP DM candidate. Assuming standard cosmology, the freeze-out abundance of the WIMP can account for the full DM abundance in part of the allowed parameter space.  It is also possible for the WIMP to form an ${\\cal O}(1)$ fraction of the DM, or a negligible fraction of DM, in other regions of the allowed parameter space.  The WIMP DM candidate in the framework is completely stable.  In addition, the tension with EDM searches is largely relaxed or inexistent.  A crucial feature of the model is the existence of an (approximate) $U(1)_{R}$-symmetry, leading to the following important consequences:  \\begin{itemize} \\item The $R$-symmetry leads to new interactions in the scalar potential which help make the phase transition strongly first-order.  \\item Since global symmetries are expected to be violated by gravitational effects, ``small\" $R$-breaking effects are included. The $R$-symmetry can thus be thought of as an accidental symmetry of the low-energy theory.  As we will see, the small $R$-breaking helps in achieving many phenomenologically desirable features at the same time:\\\\ a) {\\it Viable down-type fermion masses in some realizations (others do not require $R$-breaking)}.  \\\\ b) {\\it Evade stringent direct-detection bounds from XENON100}.  \\\\ c) {\\it Obtain a relic abundance consistent with observations in part of parameter space}.  \\item The CP-violation required for baryogenesis may have two origins: a) There are explicit CP phases in the parameters of the scalar potential, resulting in complex, spacetime-dependent \\vevs, inducing CP-violating phases in the mass matrices for the charginos and neutralinos, \\emph{or} b) All parameters in the scalar potential are real and CP violation arises from the phases in the $R$-symmetry breaking Majorana masses of the charginos and neutralinos. In both cases, the approximate $R$-symmetry suppresses the contribution to EDMs relative to that in the non $R$-symmetric case. \\end{itemize}  The plan of the paper is as follows.  In Section~\\ref{features} we present the relevant features of the $R$-symmetric scenarios, and discuss a number of realizations and theoretical considerations.  In Section~\\ref{Potential}, we discuss the zero-temperature properties of such models, while in Section~\\ref{pot-finite} we analyze the finite-temperature potential and explain the origin of the strongly first-order phase transition.  In Section~\\ref{DM}, we discuss several aspects related to DM (relic density, direct and indirect searches), and its connection to the first-order electroweak phase transition (EWPT).  In Section~\\ref{phases} we briefly comment on the generation of the baryon-asymmetry of the Universe (BAU), and discuss the reasons that allow to satisfy EDM constraints, even with sfermion masses at around a TeV. Finally, we consider other signatures in Section~\\ref{other}, and conclude in Section~\\ref{conclusions}.  We also include two appendices with technical details. ",
        "conclusions": "\\label{conclusions}  In this work we have considered a supersymmetric extension of the SM that have an approximate $U(1)_{R}$ symmetry.  We point out that this class of models can naturally lead to a strongly first-order electroweak phase transition, and possibly to the successful production of the observed baryon asymmetry, without the tensions that plague the MSSM. We have further pointed out that there is a close connection between the EWPT and the properties of DM (see~\\cite{Cirigliano:2006dg} for a study of the DM/EWBG connection in the context of the MSSM).  The basic observations derive from the (pseudo) Dirac nature of gauginos in $R$-symmetric models.  This is motivated by the fact that Majorana gaugino masses are forbidden by the $R$-symmetry, which leads one to introduce superfields in the adjoint representation of the SM gauge group.  In particular, the Dirac partner of the bino arises from a SM singlet superfield, whose scalar component plays an essential role in leading to a first-order EWPT. Since this is a tree-level effect, the phase transition is easily strong, even for Higgs masses comfortably above the Higgs LEP bound, thus creating the necessary conditions for generating the BAU. We find that the interesting region of parameter space is also characterized by a sizable Higgsino component of the lightest neutralino (the LSP), which results in an appealing DM/EWBG connection.  This fact, together with the (pseudo) Dirac nature of the LSP, leads to rather interesting DM physics: at the time of freeze-out, the LSP behaves similarly to a Dirac fermion, and has a sizable annihilation cross-section into fermion pairs via Z-exchange.  This cross-section can be modulated by a small Majorana mass splitting.  On the other hand, for annihilations today, or in scattering against nuclei, the LSP behaves like a Majorana fermion, thus being consistent with current constraints. Nevertheless, one expects that such a DM candidate will be observable in upcoming direct detection experiments, and possibly also at neutrino telescopes, such as IceCube/DeepCore.  We have further stressed that, unlike in the MSSM, the CP-violation that is relevant in the generation of the BAU is typically not constrained by EDM searches (although EDM signals are possible within the framework).  We have also commented on the collider prospects, which are characterized by a rich Higgs sector, and a possibly very interesting signal associated with the two semi-degenerate lightest neutralino states - in the form of a visible displaced vertex in association with missing energy, arising from the decay of $\\chi_2^0$ to $\\chi_1^0$.  Moreover, the collider measurements of the decay length $L$, $m_{\\chi_1^0}$ and $\\delta m$ allows the ``measurement\" of a cosmological observable - the relic abundance of the LSP. Due to the strength of the EWPT, gravitational wave signals represent another exciting possibility.  In conclusion, the present class of models represents a well-motivated extension of the SM (addressing the hierarchy problem, while significantly alleviating the SUSY flavor and CP problems), that can also easily lead to a successful generation of the BAU during the EWPT, while providing a non-standard DM candidate.  In all these regards, it compares favorably to the MSSM. More detailed studies of the points made above are worth pursuing and are left for future work."
    },
    "1107/1107.2936_arXiv.txt": {
        "abstract": "Tiny grains such as polycyclic aromatic hydrocarbons (PAHs) have been thought to dramatically reduce the coupling between gas and magnetic fields in weakly ionized gas such as in protoplanetary disks (PPDs) because they provide tremendous surface area to recombine free electrons. The presence of tiny grains in PPDs thus raises the question of whether the magnetorotational instability (MRI) is able to drive rapid accretion to be consistent with observations. Charged tiny grains have similar conduction properties as ions, whose presence leads to qualitatively new behaviors in the conductivity tensor, characterized by $\\bar{n}/n_e>1$, where $n_e$ and $\\bar{n}$ denote the number densities of free electrons and all other charged species respectively. In particular, Ohmic conductivity becomes dominated by charged grains rather than electrons when $\\bar{n}/n_e$ exceeds about $10^3$, and Hall and ambipolar diffusion (AD) coefficients are reduced by a factor of $(\\bar{n}/n_e)^2$ in the AD dominated regime relative to that in the Ohmic regime. Applying the methodology of Bai (2011), we find that in PPDs, when PAHs are sufficiently abundant ($\\gtrsim10^{-9}$ per H$_2$ molecule), there exists a transition radius $r_{\\rm trans}$ of about $10-20$ AU, beyond which the MRI active layer extends to the disk midplane. At $r<r_{\\rm trans}$, the optimistically predicted MRI-driven accretion rate $\\dot{M}$ is one to two orders of magnitude smaller than that in the grain-free case, which is too small compared with the observed rates, but is in general no smaller than the predicted $\\dot{M}$ with solar-abundance $0.1\\mu$m grains. At $r>r_{\\rm trans}$, we find that remarkably, the predicted $\\dot{M}$ exceeds the grain-free case due to a net reduction of AD by charged tiny grains, and reaches a few times $10^{-8}M_{\\bigodot}$ yr$^{-1}$. This is sufficient to account for the observed $\\dot{M}$ in transitional disks. Larger grains ($\\gtrsim0.1\\mu$m) are too massive to reach such high abundance as tiny grains and to facilitate the accretion process. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.2870_arXiv.txt": {
        "abstract": "{We present the preliminary analysis of the GRB light curves obtained by \\textit{\\textit{Swift}}/XRT between November 2004 and December 2010. ",
        "introduction": "The \\textit{Swift} satellite \\citep{swift}, launched on November 2004, during the past six years of operation detected and observed more than 600 Gamma-Ray Bursts (GRBs). The vast majority ($\\sim$67\\%) of these bursts were monitored in the soft X-ray band by the X-Ray Telescope (XRT, \\citealt{ciao}) starting as early as $\\sim$80 s after the trigger. The standard model explains the X-ray afterglow of GRBs as synchrotron radiation arising from the deceleration of a relativistic blast wave into the external medium.  The XRT follow up, therefore, is comprehensive of the tail of the prompt emission and of the afterglow. The XRT sample is now large enough to justify a statistical approach aimed at collecting, classifying and understanding the observational information of a wide and homogeneous sample of GRBs. This work will provide the most complete view of the X-ray properties of GRBs to date any existing theoretical model is asked to explain, while serving as a guide for future theoretical developments. The catalogue\\footnote{Margutti et al. in preparation}  is now under completion, here we report the status of the work. ",
        "conclusions": "The temporal and spectral properties of a homogenous sample of 437 GRBs have been extracted. For the first time, we have a large and homogeneous data set that allows us to perform a statistical study of the X-ray properties of GRB afterglow. Notably, the sample of GRBs with redshift information comprises 165 elements: for this sub-sample it is possible to study the intrinsic properties of these explosions. \\\\ From Fig. 2 and Tab. 1, we notice that the probability to observe a Type I or Type II XRT light curve is higher if the data are complete (117/153 for Type I and 135/148 for Type II) evidencing an observational bias. The majority of incomplete light curves are Type 0 (69/118). In addition, almost all the incomplete light curves have not flares (103/118) and 33\\% of the complete light curves have flares (107/319; see \\citealt{chincarini10}).\\\\ As an example of the output of our analysis we plot the Dainotti relation \\citep{dainotti08,dainotti10,dainotti11}, that is the relation between the end plateau luminosity and time, as we derive it from the row output and the distribution (histogram) of the total X-ray energy where the uncertainties for each bin have been computed by Monte Carlo simulations (Fig. 3).  \\begin{figure} \\begin{center} \\includegraphics[width=0.4\\textwidth]{End_plateau_rest_Lum_end_plateau_log_poster.eps} \\includegraphics[width=0.4\\textwidth]{Etot_X_log_GAUSS.eps} \\caption{\\small{The Dainotti relation (left) and the histrogram of the total X-ray energy with errors calculated with the Monte Carlo method (right). In red the short GRBs.}} \\end{center} \\label{esempi} \\end{figure}  \\small{"
    },
    "1107/1107.3451_arXiv.txt": {
        "abstract": "{The Solar Dynamics Observatory (SDO) was launched in February 2010 and is now providing an unprecedented view of the solar activity at high spatial resolution and high cadence covering a broad range of temperature layers of the atmosphere.} {We aim at defining the structure of a polar crown cavity and describing its evolution during the erupting process. } {We use high cadence time series of SDO/AIA observations at 304\\AA~(50000 K) and 171\\AA~(0.6 MK) to determine the structure of the polar crown cavity and its associated plasma as well as the evolution of the cavity during the different phases of the eruption. } {We observe coronal plasma shaped by magnetic field lines with a negative curvature (U-shape) sitting at the bottom of a cavity. The cavity is located just above the polar crown filament material. We thus observe the inner part of the cavity above the filament as depicted in the classical three part Coronal Mass Ejection (CME) model composed of a filament, a cavity and a CME front. The filament (in this case a polar crown filament) is part of the cavity and makes a continuous structuring from the filament to the CME front depicted by concentric ellipses (in a 2D cartoon).} {We propose to define a polar crown cavity as a density depletion sitting above denser polar crown filament plasma drained down the cavity due to gravity. As part of the polar crown filament, plasma at different temperatures (ranging from 50000K to 0.6 MK) is observed at the same location on the cavity dips and sustained by a competition between the gravity and the curvature of magnetic field lines. The eruption of the polar crown cavity as a solid body can be decomposed into two phases: a slow rise at a speed of 0.6 km$\\cdot$s$^{-1}$, and an acceleration phase at a mean speed of 25 km$\\cdot$s$^{-1}$.} ",
        "introduction": "Launched in February 2010, the Solar Dynamics Observatory (SDO) is the first NASA Living With a Star mission. SDO has on board three different instruments dedicated to study the magnetic and plasma evolution of the solar corona, its associated eruptive events and their consequences on the Earth. Here we analyse observations from the Atmospheric Imager Assembly (AIA) with high time cadence ($\\sim$ 12s) and  spatial resolution ($\\sim$ 1\\arcsec) which provide much  more detail of the corona on the limb than other comparable instruments such as SOHO/EIT and STEREO/SECCHI/EUVI. Especially it has been shown that SDO/AIA has the sensitivity to observe far more off-limb structures than never before \\citep{lem11}. Here we report on one of these structures: a cavity observed on 13 June 2010 associated with an erupting polar crown filament/prominence leading to a propagating Coronal Mass Ejection (CME) at a speed of about 300-350 km$\\cdot$s$^{-1}$ (as reported in the SOHO/LASCO CME catalog,  http://cdaw.gsfc.nasa.gov/CME\\_list/).  Typically CMEs are observed in white light coronograph with a three part structure: the bright core related to the filament material, the cavity surrounding the core and the bright front of the CME marking the transition between the CME and the ambient corona \\citep[see e.g.,][]{ill86}. In a recent paper by \\citet{gib10}, the authors studied in detail the structure, shape and  evolution (mostly due to the solar rotation) of a stable cavity observed by SOHO/EIT in the wavelength filter at 195\\AA. They found that, based on a forward modelling approach, the cavity is darker than the surrounding due to the depletion in density by a factor of about 2. It is clear in their study that the prominence cavity analysed corresponds to the classical cavity as defined in the three-part CME model mentioned above.  The formation and instability of cavities associated with a CME have been extensively modelled \\citep[see reviews by][]{lin02,for06}. The structure often contains a twisted flux tube able to accumulate magnetic energy and mass within it which thus can be destabilised by catastrophic mechanisms and/or external triggers. Despite several studies of the thermal structures of cavities especially from white-light coronographs, eclipse observations, EUV and soft  X-ray imaging  \\citep[e.g.,][]{hud99a,hud00,gib06,hab10}, there is to our knowledge no observational evidence of a long time series and high cadence obsevations of cavity at different temperatures as provided by SDO/AIA data able to (i) clearly demonstrate the thermal structure of both the prominence and cavity material, and (ii) describe how the plasma of a polar-crown filament evolves before and during the eruption.  Thus, the observations described here focus for the first time on the dynamics of the inner part of the cavity above the polar crown filament/prominence material and its evolution during the eruptive phase.  We first describe the multithermal structure of the polar crown filament/cavity (Section~\\ref{sec:aia_obs}) and then the evolution of the plasma during the eruption in two different wavelengths (Section~\\ref{sec:cavity}). In Section~\\ref{sec:disc}, we propose a definition for a polar crown cavity and discuss the implications of our study on eruptive filament models. ",
        "conclusions": "\\label{sec:disc}  We propose to define a polar crown cavity as a density depletion at the bottom of which the polar crown filament material sits indicating the existence of a magnetohydrostatic equilibrium. The filament material is drained down along the polar crown cavity by gravity and sustained by the action of the upward-directed magnetic field curvature force. This fact as well as the long, steady observations of the polar crown cavity indicate that the cavity is in a magnetohydrostatic equilibrium. The cold and hot coronal plasma are located at a similar location along the same field lines. The observations of the cavity structure and plasma spatial distribution are consistent with the classical 2D cartoon of a cavity depicted by concentric ellipses. For instance, in the classical {\\em  CSHKP} model \\citep{car64,stu68,hir74,kop76}, the eruptive structure is composed of a twisted flux tube at the bottom of which the plasma is concentrated. Contrary to the cartoon proposed by \\citet{cli86} placing the filament material at the centre of the cavity, the filament material is located at the bottom of the cavity which is more consistent with the model of \\citet{mar89}.  In the observations reported here,  magnetic curvature compensates gravity to create an equilibrium state in which the density is considerably increased at the bottom of the cavity. We also show that the flows along field lines and varying from one wavelength to an other are important for the initiation (Fig.~\\ref{fig:evol_cav}b) and relaxation (Fig.~\\ref{fig:evol_cav}d) phases of the eruption. We also note that the rise of the cavity (divided in two stages) is similar to the  plasmoid eruption initiated by an impulsive flare as reported by \\citet{ohy97} \\citep[see also][]{shi98}.  This preliminary study describes the structure and evolution of a polar crown filament and its cavity projected onto the plane of the sky and, in any case, gives a full picture of the eruption process. In a forthcoming paper, we will discuss the possible triggers of the cavity eruption of concomitant external (flares, CMEs) and internal (kink/torus instability, mass loading) phenomena by combining SDO/AIA and STEREO/SECCHI/EUVI images which give us a more realistic 3D representation of the event."
    },
    "1107/1107.3517_arXiv.txt": {
        "abstract": "{$\\epsilon$ Auriga is one of the most mysterious objects on the sky. Prior modeling of its light-curve assumed a dark, inclined, non-transparent or semi-transparent, dusty disk with a central hole. The hole was necessary to explain the light-curve with a sharp mid-eclipse brightening. } {The aim of the present paper is to study the effects of dust on the light-curves of eclipsing binary stars and to develop an alternative physical model for $\\epsilon$ Aur type objects which is based on the optical properties of dust grains. } {The code Shellspec has been modified to calculate the light-curves and spectra of such objects. The code solves the radiative transfer along the line of sight in interacting binaries. Dust and angle dependent Mie scattering were introduced into the code for this purpose. } {Our model of $\\epsilon$ Aur consists of two geometrically thick flared disks: an internal optically thick disk and an external optically thin disk which absorbs and scatters radiation. Disks are in the orbital plane and are almost edge-on. We argue that there is no need for an inclined disk with a hole to explain the current eclipse of $\\epsilon$ Aur not even if there is a possible shallow mid-eclipse brightening. It was demonstrated that phase dependent light scattering and the optical properties of the dust can have an important effect on the light-curves of such stars and can even produce a mid-eclipse brightening. This is a natural consequence of the strong forward scattering. It was also demonstrated that shallow mid-eclipse brightening might result from eclipses by nearly edge-on flared (dusty or gaseous) disks. } {} ",
        "introduction": "$\\epsilon$ Aur has eluded astronomers for a long time. This star is a bright object visible by the naked eye and known to be variable for almost two hundred years. It is an eclipsing binary with the longest known orbital period, 27.1 yr. A very rare eclipse which lasts for 2 years is almost over. However, this binary star, its origin and, in particular, its secondary component, remain mysterious. The extremely long orbital period and eclipse implies that the two objects are huge.  The primary, main source of light, may be a relatively young, massive F0Ia super-giant with the mass of about 16 $M_{\\odot}$ (Stencel et al. \\cite{stencel}) and radius of about 135 $R_{\\odot}$ (Hoard et al. \\cite{hoard}). The secondary is unseen. Kuiper, Struve \\& Str\\\"{o}mgren (\\cite{kuiper}) suggested that it is a huge ($>3000R_{\\odot}$) partially transparent star. Nowadays, we believe that the secondary is a dark and mysterious object with a radius of about 9 AU and a mass of 13 $M_{\\odot}$. Huang (\\cite{huang}) proposed that the secondary is a dark disk seen edge on. Wilson (\\cite{wilson}) and Carroll et al. (\\cite{carroll}) argue that the observed sharp mid-eclipse brightening {\\bf (MEB)} can only be explained by a tilted disk with a central opening. Ferluga (\\cite{ferluga}) suggested that the disk is a system of rings. \\footnote{By a mid-eclipse brightening in $\\epsilon$ Aur some authors understand only a relatively sharp local maximum which appeared in a few eclipses near the middle of the eclipse. We propose a slightly more general definition of the MEB. Our MEB or Mid-Eclipse Excess is a convex feature near the middle of an eclipse bed. A common eclipse has a concave eclipse bed. In the pictures with reversed magnitudes it looks just the opposite.} The main weakness of this model is that it cannot explain what is inside the disk and why we do not see it. Stars with enough mass are too bright and black holes are ruled out by the lack of X-rays. An alternative is that the primary is an old evolved post-AGB star with a mass of about 2.2 $M_{\\odot}$ and that the secondary hides a main sequence B5 star with mass of about 5.9 $M_{\\odot}$ in the center (Hoard et al. \\cite{hoard}, Kloppenborg et al. \\cite{kloppenborg2010}, Takeuti \\cite{takeuti} and references therein).  There has been a wealth of studies during the current eclipse at various wavelengths. Orbital solutions were recently revisited by Stefanik et al. (\\cite{stefanik}) and Chadima et al. (\\cite{chadima2010}). Kloppenborg et al. (\\cite{kloppenborg2010}, \\cite{kloppenborg2011}) confirmed the dark disk using interferometric observations but they did not confirm the hole in the disk. The spectral energy distribution was studied by Hoard et al. (\\cite{hoard}). They favor the post AGB+B5V model and concluded that the dust grains in the disk are unusually large (see also Lissauer et al. \\cite{lissauer}). Wolk et al. (\\cite{wolk}) analyzed X-ray observations. Recently, Chadima et al. (\\cite{chadima2010}, \\cite{chadima2011b}) questioned the presence of sharp mid-eclipse brightening and suggested that the photometric variability seen during eclipse is intrinsic to the F star.  The section below presents a motivation for this study and an outline of important physical and geometrical effects. These effects, dust and Mie scattering were incorporated into the code Shellspec. It solves radiative transfer along the line of sight in 3D moving circumstellar environment (Budaj \\& Richards \\cite{budaj2004}, see also Budaj \\cite{budaj2011}, Miller et al. \\cite{miller}, Tkachenko et al. \\cite{tkachenko}, Chadima et al. \\cite{chadima2011a} for some modifications and applications). In the next sections we apply this code to $\\epsilon$ Aur. ",
        "conclusions": "A new alternative model of the $\\epsilon$ Aur type stars was suggested. This model is based on the optical properties of dust grains. It assumes that dust concentrates in two regions: an optically thick disk and optically thin regions of different shapes. It was demonstrated that an optically thick flared disk coated with an optically thin layer can reproduce the current eclipse observations of $\\epsilon$ Aur very well. The model takes into account Mie scattering on dust, i.e. extinction due to the absorption and scattering, as well as thermal and scattering emission.  It was demonstrated that angle dependent Mie scattering is important in such systems. Under certain circumstances, mainly if the dark dust is eclipsing a strong source of light, the forward scattering on dust can be crucial and can even produce a mid-eclipse brightening. It was also demonstrated that near edge-on flared disks can produce eclipses with a shallow mid-eclipse brightening and that this might be an explanation of the shallow mid-eclipse brightening feature of $\\epsilon$ Aur.  This model provides an alternative to an inclined disk with a central hole invoked to explain systems with mid-eclipse brightening. It does not raise the question why we do not see the object in the center of the disk. We may not see it because in our model it may be obscured by the disk. Potential variability in the mid-eclipse brightening might be caused by subtle changes in the remote, low density, optically thin regions of dust located at small gravitational potential, or by small changes in the inclination (or warping) of the disk and does not have to involve dense matter at high gravitational potential in the central hole of the disk."
    },
    "1107/1107.4337_arXiv.txt": {
        "abstract": "In 4-dimensional General Relativity, there are several theorems restricting the topology of the event horizon of a black hole. In the stationary case, black holes must have a spherical horizon, while a toroidal spatial topology is allowed only for a short time. In this paper, we consider spinning black holes inspired by Loop Quantum Gravity and by alternative theories of gravity. We show that the spatial topology of the event horizon of these objects changes when the spin parameter exceeds a critical value and we argue that the phenomenon may be quite common for non-Kerr black holes. Such a possibility may be relevant in astrophysics, as in some models the accretion process can induce the topology transition of the horizon. ",
        "introduction": "The study of black hole (BH) uniqueness theorems started more than forty years ago and it is still a very active research field~\\cite{review}. In 4-dimensional General Relativity, the Hawking's theorem ensures that the spatial topology of the event horizon must be a 2-sphere in the stationary case, under the main assumptions of asymptotically flat space-time and validity of the dominant energy condition~\\cite{hawking}. Technically, the event horizon of a BH is defined as the boundary of the causal past of future null infinity. The spatial topology of the event horizon at a given time is the intersection of the Cauchy hypersurface at that time with the event horizon. According to the topological censorship theorem, in a globally hyperbolic and asymptotically flat space-time, any two causal curves extending from past to future infinity are homotopic~\\cite{tct}. As a BH with a toroidal spatial topology would violate this theorem, the hole must quickly close up, before a light ray can pass through~\\cite{tedj}. Interestingly, numerical simulations find that toroidal horizons can form, but they exist for a short time, consistently with the topological censorship theorem~\\cite{teuk}.   In the vacuum, the only stationary and axisymmetric BH solution of the Einstein's equations in a 4-dimensional and asymptotically flat space-time is given by the Kerr geometry, which is completely specified by two parameters: the mass $M$ and the spin angular momentum $J$ -- instead of $J$, it is more commonly used the spin parameter $a = J/M$ or the dimensionless spin parameters $a_* = J/M^2$. While the Kerr metric may be a solution even in other theories of gravity, in general there is not a similar uniqueness theorem~\\cite{kerr}. Unfortunately, for the time being our knowledge of non-Kerr BH solutions in alternative theories of gravity is definitively limited. In most cases, analytic or numerical metrics are known only for non-rotating BHs. Approximated solutions in the slow-rotation limit have been obtained in Einstein-Gauss-Bonnet-dilaton (EGBd) gravity~\\cite{slow1} and in Chern-Simons modified gravity~\\cite{slow2}. The only 4-dimensional non-Kerr spinning BH example of a specific gravity theory is given by a numerical metric in EGBd gravity, recently found in~\\cite{egbg}. On the other hand, from the observational point of view, fast-rotating BHs are the most interesting, as deviations from the Kerr metric are typically more evident and we would have more chances to test the model with astrophysical data~\\cite{rev-b}.   In this paper, we discuss two examples of spinning BHs in theories beyond General Relativity, proposed respectively in~\\cite{bh2} and \\cite{dim}. In both cases, the two metrics have not been obtained by solving specific field equations, but the fact they have an analytic form for arbitrary values of the spin parameter is very useful. The BH proposed in~\\cite{bh2} is inspired by Loop Quantum Gravity and the deformations from the Kerr geometry are encoded in a polymeric function $P$ and in a Plank scale parameter $a_0$. The BH proposed in~\\cite{dim} is instead a phenomenological metric and may be seen as a BH solution in an unknown alternative theory of gravity\\footnote{These metrics can be solutions of particular non-local generalizations of the Einstein's equations, as suggested in~\\cite{ModestoMoffatNico,SRQG}. The idea is to replace the Einstein's equations with the following set of equations of motion $$  R_{\\mu \\nu} -\\frac{1}{2} g_{\\mu \\nu} R = 8 \\pi G_N \\mathcal{O}( \\Box/\\Lambda^2) T_{\\mu \\nu} \\, ,$$ where $\\mathcal{O}(\\Box/\\Lambda^2)$ is a generic non-local function of the covariant D'Alembertian operator and $\\Lambda$ is the energy scale of the modified gravity~\\cite{SRQG}.  }. Both BHs have been obtained through a Newman-Janis transformation of a static solution. The key point of the two metrics is that there are no restrictions on the values of the spin parameter $a_*$. Interestingly, for high values of $a_*$ they present similar features. If the BH is more prolate than the predictions of General Relativity, when it rotates fast the spatial topology of the event horizon changes from a 2-sphere to two disconnected 2-spheres. If the BH is more oblate, one finds a toroidal horizon or something very similar to a toroidal horizon. Our guess is that fast-rotating BHs with non-trivial topology are not peculiar predictions of these two solutions, but that they may be relatively common in the case of deviations from the Kerr geometry. ",
        "conclusions": "In 4-dimensional General Relativity, a stationary BH must have a spherical horizon, while a toroidal horizon is allowed for a very short time. It is thus thought that the spatial topology of the horizon of astrophysical BHs is a 2-sphere. However, if the current BHs candidates are not the BHs predicted by General Relativity, this conclusion may be wrong. Here we have discussed two examples of 4-dimensional non-Kerr spinning BHs and we have shown that for high values of the spin parameter the topology of the event horizon of these objects can change. Unfortunately, our current knowledge of these objects is definitively limited. Here we have considered the loop-improved Kerr metric found in Ref.~\\cite{bh2} and the phenomenological metric proposed in~\\cite{dim} to perform tests of strong gravity. Interestingly, they present quite remarkable and qualitatively similar features in the case of fast-rotating objects and we guess that these properties may be common for non-Kerr spinning BHs. In particular, it seems that BHs more prolate than the Kerr one develop two disconnected topologically spherical horizons above some critical spin parameter. In the case of fast-rotating objects more oblate than a Kerr BH, their horizon looks more like a torus, even if the central hole may be closed (like in the case of the BH in the right panel of Fig.~\\ref{f2}, in which the event horizon extends up to the central singularity at $r=0$). In both this examples, the accretion process may overspin these objects above the critical spin parameter and induce the topology transition of the horizon. The topology change does not happen as a result of some jump or tunneling, but this fact should not be seen suspiciously: even in numerical simulations in General Relativity the gravitational collapse can produce a toroidal BH and then the hole quickly closes up continuously, as found for the first time in~\\cite{teuk}.   As final remark, let us notice that here we have not discussed the stability of these BHs. However, this issue cannot be addressed for the metrics~(\\ref{eq-metric}) and (\\ref{eq-ps}): we do not know the field equations of the gravity theory having Eqs.~(\\ref{eq-metric}) and (\\ref{eq-ps}) as solution, and therefore we cannot predict the evolution of small perturbations on these backgrounds. On very general grounds, we can simply say the event horizon of very fast-rotating objects becomes likely too small to prevent the ergoregion instability~\\cite{ergo}. However, such an instability may occur only for BHs with spin parameters $a_* > a_*^{eq}$, which would be anyway unstable configurations. If, on the contrary, the instability appears at $a_* < a_*^{eq}$, as the spin-up due to the accretion process is an unavoidable phenomenon, these BHs would be a source of gravitational waves, potentially detectable by future experiments."
    },
    "1107/1107.3721_arXiv.txt": {
        "abstract": "We propose to use the threshold-free process of neutrino capture on $\\beta$-decaying nuclei (NCB) using all available candidate nuclei in the Milky Way as target material in order to detect the presence of the Cosmic neutrino background (C$\\nu$B). By integrating over the lifetime of the galaxy one might be able to see the effect of NCB processes as a slightly eschewed abundance ratio of selected $\\beta$-decaying nuclei. First, the candidates must be chosen so that both the mother and daughter nuclei have a lifetime comparable to that of the Milky Way or the signal could be easily washed out by additional decays. Secondly, relic neutrinos have so low energy that their de Broglie wavelengths are macroscopic and they may therefore scatter coherently on the electronic cloud of the candidate atoms. One must therefore compare the cross sections for the two processes (induced $\\beta$-decay by neutrino capture, and coherent scattering of the neutrinos on atomic nuclei) before drawing any conclusions. Finally, the density of target nuclei in the galaxy must be calculated. We assume supernovae as the only production source and approximate the neutrino density as a homogenous background. Here we perform the full calculation for $^{187}$Re and $^{138}$La and find that one needs abundance measurements with 24 digit precision in order to detect the effect of relic neutrinos. Or alternatively an enhancement of $\\rho_{\\nu}$ by a factor of $\\sim10^{15}$ to produce an effect within the current abundance measurement precision. ",
        "introduction": "In analogy with the photons of the Cosmic Microwave Background (CMB), the neutrino component of the early Universe is expected to have decoupled from the hot primordial plasma once the rate of interactions with the rest of the particles dropped below the Hubble expansion rate \\cite{Fuku:03}. The neutrinos travelling towards us from this last scattering surface are normally referred to as the Cosmic Neutrino Background (C$\\nu$B) or relic neutrinos.  A detection of the C$\\nu$B would be of great interest to cosmology as well as neutrino physics. Contrary to other neutrino sources outside the solar system, the C$\\nu$B has the advantage of being a homogenous and isotropic background whose number density is only superseded by the CMB photon background. The theoretical expected number density is $n_{\\nu + \\overline{\\nu}}$=112 cm$^{-3}$ for each neutrino mass state \\cite{Fuku:03}. However, this advantage pales somewhat in comparison with the fact that the temperature of the background neutrinos is so low that their mean energy is many orders of magnitude below the threshold of any of the classical neutrino detection methods.  One can calculate the C$\\nu$B temperature, seeing as it is closely related to the very accurately measured photon temperature, $T_{\\gamma}=2.72548\\pm0.00057$ K \\cite{Fixsen:09}. Using entropy conservation, the photon temperature before reheating -- by electron-positron annihilation -- can be related to the temperature after reheating. The entropy is proportional to the number of degrees of freedom times the temperature cubed and therefore the final temperature is given by: \\begin{equation} \\centering g_{*,f}T_f^3=g_{*,i}T_i^3. \\end{equation} The number of degrees of freedom before reheating comes from both photons and electrons, so that $g_{*,i}=1+7/4$. After reheating the only source of radiation is photons, i.e.~$g_{*,f}=1$. By assuming the photon temperature before reheating to be equal to the neutrino temperature one gets: \\begin{equation} \\centering T_{\\nu}=\\left(\\frac{4}{11}\\right)^{1/3} T_{\\gamma}. \\end{equation} Consequently the neutrino temperature today corresponds to a neutrino energy of only 6.5$T_{\\nu}^2/m_{\\nu}$ or 3.15$T_{\\nu}$ for non-relativistic and relativistic neutrinos respectively, and a temperature of $T_{\\nu}=1.95$K translates to an energy of $T_{\\nu}k_B\\sim$ 1.6 $\\cdot 10^{-4}$eV\\cite{Cocco:07}. This energy is far below the threshold for both the radio-chemical and Cherenkov neutrino detectors that have previously been used in experiments with solar, atmospheric, reactor, accelerator, and supernova neutrinos. However, it has long been known that neutrino capture on $\\beta$-decaying nuclei (NCB) does not have a threshold (if the normal $\\beta$-decay process is already spontaneously occurring in nature). Due to energy conservation, the $\\beta$-particle will receive all the energy of the decay plus one neutrino rest-mass\\footnote{Under the reasonable assumption that the daughter nucleus is so heavy that its recoil energy is negligible.}. This results in a small peak in the $\\beta$-spectrum one neutrino mass above the theoretical endpoint, which is basically a smoking gun signature for the relic neutrino background unless non-standard model physics is introduced.  The size of the peak is a measure of the number of capture processes, which depends of course on the cross section of the reaction, but also on the C$\\nu$B number density. Recent calculations show that even if one can, in principle, resolve the neutrino mass in the $\\beta$-spectrum\\footnote{The calculation in question was performed for the KArlsruhe TRItium Neutrino (KATRIN) experiment, which is expected to be the first in a new generation of $\\beta$-decay experiments able to determine the neutrino mass with sub-eV sensitivity \\cite{KAT:04}.}, the background density will have to be enhanced by roughly a factor $10^9$ in order to produce a detectable CNB signal \\cite{Kaboth:10}. Or alternatively, the number density of the source material should be enhanced by the same factor. ",
        "conclusions": "As shown in Figure~\\ref{fig:abbus} the abundance enhancement is tiny and far below the precision of modern galactic abundance measurements even when using meteoritic samples. Our results also show that one must have a background enhancement of around $10^{15}$ to see the effect even on the 8th decimal which is just around the precision available on the Osmium abundance -- see e.g.~\\cite{Ackena:11, Lodders:03}  So despite our use of an extremely large detector volume, we must unfortunately conclude that a detection is even more orders of magnitude away with this method, than say, with a $\\beta$-decay experiment such as KATRIN \\cite{KAT:04} or MARE \\cite{Monfardi:2006,Nucci:2010}."
    },
    "1107/1107.0722_arXiv.txt": {
        "abstract": "We investigate the physical processes involved in the development of the red sequence (RS) of cluster galaxies by using a combination of cosmological {\\em N}-body simulations of clusters of galaxies and a semi-analytic model of galaxy formation. Results show good agreement between the general trend of the simulated RS and the observed colour-magnitude relation (CMR) of early-type galaxies in different magnitude planes. However, in many clusters, the most luminous galaxies ($M_R \\sim M_V \\sim M_{T_1} \\lesssim -20$) depart from the linear fit to observed data, as traced by less luminous ones, displaying almost constant colours.  With the aim of understanding this particular behaviour of galaxies in the bright end of the RS, we analyze the dependence with redshift of the fraction of stellar mass contributed to each galaxy by different processes, i.e., quiescent star formation, and starbursts triggered by disc instability and merger events. We find that the evolution of galaxies in the bright end since $z\\approx 2$ is mainly driven by minor and major dry mergers, while minor and major wet mergers are relevant in determining the properties of less luminous galaxies. Since the most luminous galaxies have a narrow spread in ages ($1.0\\times 10^{10}$ yr $<t<1.2\\times 10^{10}$ yr), their metallicities are the main factor that affects their colours.  Their mean iron abundances are close to the solar value and have already been reached at $z \\approx 1$. This fact is consistent with several observational evidences that favour a scenario in which both the slope and scatter of the CMR are in place since $z \\approx 1.2$. Galaxies in the bright end reach an upper limit in metallicity as a result of the competition of the mass of stars and metals provided by the star formation occuring in the galaxies themselves and by the accretion of merging satellites.  Star formation activity in massive galaxies (stellar mass $M_\\star \\gtrsim 10^{10} M_{\\odot}$) that takes place at low redshifts contribute with stellar components of high metallicity, but the fraction of stellar mass contributed since $z\\approx 1$ is negligible with respect to the total mass of the galaxy at $z=0$.  On the other hand, mergers contribute with a larger fraction of stellar mass ($\\approx 10-20$ per cent), but the metallicity of the accreted satellites is lower by $\\approx 0.2$ dex than the mean metallicity of galaxies they merge with. The effect of dry mergers is to increase the mass of galaxies in the bright end, without significantly altering their metallicities. Hence, very luminous galaxies present similar colours that are bluer than those expected if recent star formation activity were higher, thus giving rise to a break in the RS.  These results are found for simulated clusters with different virial masses ($10^{14} - 10^{15}\\,h^{-1}\\,{\\rm M}_\\odot$), supporting the idea of the universality of the CMR in agreement with observational results. ",
        "introduction": "It is now well established that galaxies follow a bimodal distribution in the colour-magnitude plane, separated into a tight `red sequence' (RS) and a `blue cloud'. The RS mostly consists of gas-poor galaxies with low levels of star formation (SF), prototypically early-type galaxies (ETGs), whereas late-type galaxies are typical objects of the blue cloud.  The colour-magnitude relation (CMR) of ETGs can be understood as a mass-metallicity relation. The more luminous (and therefore more massive) galaxies in this relation have deep potential wells, capable of retaining the metals released by supernovae events and stellar winds. The CMR seems to be quite universal, since it is followed by galaxies in the field as well as in groups and clusters \\citep[e.g.][]{Visvanathan77,Sandage78,Reda04,LopezCruz04,Reda05,deRijcke09}, but with a larger fraction of red galaxies being in denser environments.  There is controversy on which are the mechanisms responsible for the evolution of galaxies to the RS. The physical processes invoked can be classified in internal and external. The former include passive stellar evolution and disc instabilities, while the latter involve galaxy-galaxy interactions and mergers \\citep{Lin10,Robaina10}, and environmental effects such as fast gravitational interactions (`galaxy harassment', \\citealt{Moore99}), removal of the hot gas reservoir (`strangulation' or `starvation', \\citealt{Larson80,Balogh00,Kawata08}) and ram pressure stripping of galactic gas \\citep{Lanzoni05, Roediger09,Tecce10}.  Studying the evolution of the slope and scatter of the CMR provides constraints on the SF history of ETGs. \\citet{Bower98} find that the observed CMR scatter in local clusters can be reproduced if galaxies are assumed to form the bulk of their stellar content at early epochs ($z > 1$), with a subsequent moderate mass growth driven by merger events and residual SF activity. This scenario is consistent with hierarchical models of galaxy formation. In this context, using semi-analytical modelling, \\citet{Menci08} obtain a narrow RS for cluster galaxies, already defined by $z \\approx 1.2$. At a similar redshift, \\citet{Kaviraj05} find that the slope of the CMR of cluster galaxies changes appreciably with respect to its present value, although the evolution of the slope in the redshift range $0 < z < 0.8$ is negligible taking into account the errors.  Detailed observations of the CMR of cluster galaxies (see \\citealt{Mei09}, and references therein) show no significant evolution of the zero point, slope and scatter of the CMR out to $z\\approx 1.3$. In contrast with these results, \\citet{Stott09} find from optical and near-infrared observations of cluster galaxies an evolution of the CMR slope between $z \\approx 0.5$ and the present epoch, which they attribute to the infall of galaxies into the cluster core that are being transformed into RS galaxies.  Hydrodynamical simulations have also been used to analyse the properties of the CMR at $z=0$. \\citet{Saro06} evaluate the impact of the stellar initial mass function (IMF), finding that the Salpeter IMF allows to recover both the slope and the normalization of the CMR for galaxies in cluster-sized haloes. From the analysis of simulations of galaxy groups and clusters, \\citet{Romeo08} conclude that the shape of the RS is mainly determined by the specific SF in all environments. These authors find that evolving galaxies move towards a `dead' sequence soon after they have almost stopped the bulk of their SF. Fainter galaxies in clusters keep a significant amount of SF out to very recent epochs, and are thus distributed more broadly around the RS.  A linear relation has generally been used to fit the correlation between luminosity and colour of the CMR, from giant to dwarf ellipticals. However, the scatter of these relations increases at lower luminosities, as becomes evident from observations of several clusters (Coma, \\citealt{Secker97}; Perseus, \\citealt{Conselice03}; Fornax, \\citealt{Hilker03}, \\citealt{Karick03}, \\citealt{Mieske07}; Hydra, \\citealt{Misgeld08}; Virgo, \\citealt{Lisker08}; Antlia, \\citealt{Analia08}).  Some of these relations seem to be consistent with a change of slope from the bright to the faint end. As an example, this kind of fit has already been suggested by \\citet{Vacu61} for the Virgo cluster, and confirmed later by \\citet[e.g.][]{LaFerrarese06}. More recently, based on Sloan Digital Sky Survey (SDSS) imaging data, \\citet{Janz09} found, also for Virgo, a non-linear relation that can be described by a `S' shape over the whole range of magnitudes, with the brightest galaxies ($-21 \\lesssim M_{\\rm B} \\lesssim -19)$ having an almost constant colour. In the case of the Hydra cluster, \\citet{Misgeld08} fit a linear relation to its CMR, but a change of slope is evident for the brightest galaxies in that cluster.  Additional evidence of a tilt towards bluer colours at the bright end of the RS arises from studies of large samples of galaxies in the SDSS \\citep[e.g.][]{Baldry04,Baldry06,Skelton09}. The results of \\citet{Skelton09} were obtained for the local RS averaged over all environments. \\citet{Baldry06} examined the dependence of the fit (using a `S'-shaped combination of straight line plus a hyperbolic tangent function) to the mean positions of both the RS and blue cloud, and found that unlike the fractions of red galaxies, the fits to the RS do not vary strongly with environment.  The detachment of the bright end from the general trend denoted by the linear fit to the CMR of cluster galaxies motivates our study. The CMR constitutes one of the major tools to test galaxy formation models since  galactic evolution is evidenced by it. Our aim is to contribute to the understanding of the physical processes behind the development of the RS in galaxy clusters, and the special behaviour of its bright end.  Recent observations support strong evolution with cosmic time in the structural properties of the most massive (stellar mass $M_\\star \\gtrsim 10^{11} M_\\odot$) spheroid-like galaxies. Size measurements of ellipticals at high redshift ($z \\gtrsim 1.5$) reveal that these objects are much smaller (by factors $\\sim 2$ to $\\sim 4$) than their local counterparts of comparable stellar mass (\\citealt{Daddi05}). This result was subsequently confirmed by several studies, although the most massive ETGs ($M_\\star \\gtrsim 2.5\\times10^{11} M_\\odot$) seem to have reached their internal and dynamical structure earlier and faster than lower mass ETGs at the same redshift, as discussed by \\citet{Mancini10}.  On the other hand, a mild evolution in velocity dispersion is found for massive galaxies \\citep{CenarroTrujillo09}.  On average, the velocity dispersion and surface mass density of massive galaxies at high redshift are similar to those of the most dense local ETG \\citep{Cappellari09}.  These observational evidences suggest that inside-out growth scenarios are plausible, in which the compact high redshift galaxies make up the centres of normal nearby ellipticals \\citep{vanDokkum10}.  From samples of galaxy pairs obtained from different surveys, it is found that present-day massive ETGs undergo, on average, $\\sim 0.5$ major mergers since $z \\sim 0.6 - 0.7$ \\citep{Bell06b,Robaina10}. These mergers generally occur between galaxies which are already on the RS, thus involving ETGs containing little cold gas and dust \\citep{WhitakervanDokkum08,McIntosh08}. These dissipationless (gas-poor, thus called `dry') major mergers are an important channel for the formation and evolution of the brightest cluster galaxies (BCGs) \\citep{Liu09}. However, the rate of dry major mergers is too low to fully explain the formation of spheroidal (and RS) galaxies since $z\\sim 1$ \\citep{Bundy09}.  Minor dry mergers would favour the growth of BCGs as inferred from the evolution of their velocity dispersion-stellar mass relation \\citep{Bernardi09}. The minor merger hypothesis also appears to be plausible for early-type galaxies at intermediate redshifts ($0.1<z<0.8$) as shown by \\citet{Nierenberg11} who find that ETGs are typically surrounded by several satellite galaxies.  Minor mergers are also probably responsible for the size evolution of massive ETGs \\citep{BezansonvanDokkum10}.  Besides, \\citet{Kaviraj11} demonstrate that the low-level of star formation activity in the ETG population at intermediate redshift is likely to be driven by minor mergers.  The influence of mergers on the evolution of ETGs inferred from observational data is in agreement with predictions obtained from models of galaxy formation. Using semi-analytic modelling techniques, \\citet{KhochfarBurkert03} find that the last major merger of present-day bright elliptical galaxies ($M_B \\lesssim -21$) occurs preferentially between bulge-dominated galaxies. Major mergers are responsible for doubling the stellar mass at the massive end of the RS since $z\\sim 1$ \\citep{Moral10}.  On the other hand, \\citet{Bournaud07} study the effect of multiple minor mergers on ETG evolution by using a {\\em N}-body simulation finding that repeated minor mergers can form elliptical galaxies without major mergers events, which are less frequent than minor mergers at moderate redshift. Remnants of these multiple mergers develop global properties that depend more on the total mass accreted during mergers than on the mass ratio of each merger.  This result is also supported by hydrodynamic simulations \\citep[e.g.][]{Naab07,Naab09} which suggest that the growth of massive galaxies ($M_\\star \\gtrsim 10^{11} M_\\odot$) by a large series of minor mergers becomes more important than the growth via major mergers; minor mergers may thus be the main driver for the late evolution of sizes and densities of ETGs. Another possible physical mechanism responsible for the size evolution of ETGs is the expulsion of a substantial fraction of the initial baryons, still in gaseous form, by quasar activity \\citep{Fan10}.  From semi-analytic modelling, \\citet{Khochfar06b} find that massive galaxies at high redshifts form in gas-rich mergers, whereas galaxies of the same mass at low redshifts form from gas-poor mergers. These considerations allow them to reproduce the observed size-redshift evolution of massive spheroidal objects. Dissipation associated with large gas fractions is the most important factor determining spheroid sizes, as found by \\citet{Hopkins09} using high resolution hydrodynamic simulations. Complementarily, also using semi-analytic models, \\citet{DeLucia07} find that later ($z < 0.5$) minor dry mergers play an important role in the mass assembly of BCGs, which acquire half of their mass via accretion of smaller galaxies.  It is clear that different types of mergers have a strong impact on the evolution of ETGs, affecting their morphological and kinematical properties, as well as their star formation history. As has been described, many works have been devoted to study these aspects but only a few have analysed in detail the role of mergers in the development of the CMR \\citep[e.g.][]{Skelton09, Bernardi10}. In particular, by using a toy model combined with galaxy merger trees extracted from a semi-analytic model, \\citet{Skelton09} find that the tilt towards bluer colours of the bright end of the RS  of local early-type galaxies can be explained by the higher fraction of dry mergers suffered by bright galaxies with respect to fainter ones. Galaxy colour is not expected to change during dry mergers, since there is no associated SF. Galaxies then move along the CMR as the mass of the system increases, with their colour remaining fixed \\citep{Bernardi07}.  In this work, we investigate this issue in detail using the semi-analytic model of galaxy formation and evolution \\sag~\\citep*[acronym for Semi-Analytic Galaxies;][hereafter LCP08]{Lagos08}, which is combined with hydrodynamical cosmological simulations of galaxy clusters by \\citet{Dolag05}. We focus our study on galaxies lying on the RS, which are selected by a colour criterion. The samples thus obtained mostly include early type galaxies, making our conclusions valid for the understanding of the development of the CMR of ETGs. Thus, we hereafter use the term RS to refer to colour-selected galaxies from our simulations, and the term CMR to refer to observed relations for ETGs.  The tilt in the bright end of the RS emerges naturally from the model once it has been calibrated to recover numerous observed galaxy properties, both locally and at high redshift. We quantify the contribution to the mass and metal content of galaxies located in the RS given by quiescent SF from the cold gas available in the galaxy disc, starbursts triggered by disc instability and merger events, and the stellar mass accreted from satellite galaxies during mergers, which are identified as minor/major and wet/dry ones.  This paper is organised as follows. In Section~\\ref{Sec:model}, we briefly describe the semi-analytic model used in this work and give details of the cluster simulations. Section~\\ref{Sec:RedSeq} presents the comparison between the simulated RS and observed CMR followed by early-type cluster galaxies in different magnitude planes, and the analysis of the luminosity-metallicity relation followed by galaxies in the RS. Section~\\ref{Sec:Proc} describes the formalism used to quantify the contribution of different processes to the mass and metallicity of galaxies; the corresponding results are presented and discussed.  Finally, in Section~\\ref{Sec:Conclu} we summarize our main findings. ",
        "conclusions": ""
    },
    "1107/1107.2511_arXiv.txt": {
        "abstract": "We present deep, high-resolution imaging of the nearby Fanaroff-Riley Class\\,I (FR\\,I) radio galaxies NGC\\,193, B2\\,0206+35, B2\\,0755+37 and M\\,84 at frequencies of 4.9 and 1.4\\,GHz using new and archival multi-configuration observations from the Very Large Array. In addition, we describe lower-resolution observations of B2\\,0326+39 and a reanalysis of our published images of 3C\\,296. All of these radio galaxies show twin jets and well-defined lobes or bridges of emission, and we examine the common properties of this class of source. We show detailed images of total intensity, brightness gradient, spectral index, degree of polarization and projected magnetic-field direction. The jet bases are very similar to those in tailed twin-jet sources and show the characteristics of decelerating, relativistic flows. Except on one side of M\\,84, we find that the jets can be traced at least as far as the ends of the lobes, where they often form structures which we call ``caps'' with sharp outer brightness gradients. Continuing, but less well collimated flows back into the lobes from the caps can often be identified by their relatively flat spectral indices.  The lobes in these radio galaxies are similar in morphology, spectral-index distribution and magnetic-field structure to those in more powerful (FR\\,II) sources, but lack hot-spots or other evidence for strong shocks at the ends of the jets. M\\,84 may be an intermediate case between lobed and tailed sources, in which one jet does not reach the end of its lobe, but disrupts to form a ``bubble''. ",
        "introduction": "\\label{Introduction}  Relativistic jets are the primary channel of energy loss from accreting supermassive black holes in many radio galaxies. They also have a major impact on their surroundings and act as accelerators of the most energetic photons (and perhaps hadrons) we observe. The present paper forms part of a study of jet physics in nearby, low-luminosity radio galaxies, specifically those with FR\\,I morphology \\citep{FR74}.  We have developed a sophisticated model of FR\\,I jets as relativistic, symmetrical, axisymmetric flows.  By fitting to deep, high-resolution radio images in total intensity and linear polarization, we have determined the three-dimensional variations of velocity, emissivity and magnetic-field ordering in five sources \\citep{LB02a,CL,CLBC,LCBH06}.  We have shown that FR\\,I jets decelerate from relativistic ($\\beta = v/c \\approx 0.8$) to sub-relativistic speeds on scales of a few kpc and that they are faster on-axis than at their edges, as expected if they entrain external material.  The physics of boundary-layer entrainment must depend on the composition and density of the surrounding medium, and in particular on whether the jets propagate in direct contact with the intergalactic medium (IGM) or are surrounded by lobes consisting primarily of tenuous and at least partially relativistic plasma. Of the five sources we have modelled, three have {\\em plumed} or {\\em tailed} outer structures wherein most of the extended emission appears to lie further from the active nucleus than the narrower jets: 3C\\,31 \\citep{LB02a,3c31ls}, B2\\,1553+24 \\citep{CL,Young} and NGC\\,315 \\citep{CLBC,ngc315ls}. We presume that their jets are in direct contact with the IGM. On the other hand, 3C\\,296 \\citep{LCBH06} has two {\\em lobes} with well-defined outer boundaries and a diffuse {\\em bridge} of emission around the jets (at least in projection).  A further source, B2\\,0326+39 \\citep{CL}, clearly has lobes, but it was unclear from published observations \\citep{Bridle91} whether these lobes surround the inner jets.  Our best-fit model for the jets in the lobed FR\\,I source 3C\\,296 \\citep{LCBH06} is unusual in that it shows a very large transverse velocity gradient across the jet except very close to the nucleus, the ratio of edge to central velocity in this jet being $\\la 0.1$, compared with values $\\approx$0.4 for the jets in three tailed sources (the value for 0326+39 is poorly determined). It is therefore of interest to examine whether jets in other FR\\,I sources whose lobes entirely surround them appear to decelerate differently with distance from the nucleus than those in tailed FR\\,I sources, or those in sources whose lobes may not extend all the way back to the nucleus, leaving the inner jets unshielded. Projection complicates interpretation of individual sources (the lobes may appear superimposed on the jets even if they are not in physical contact) and it is not always straightforward to separate jet and lobe emission, but the differences between 3C\\,296 and the rest are large. Modelling of the jets in a small number of lobed sources (which form the majority of complete samples of low-luminosity radio galaxies; \\citealt{PDF96}) should therefore be enough to decide whether there are systematic differences between the two classes.  The primary aim of this paper is to present high-quality radio imaging of four lobed FR\\,I sources whose jets are suitable for modelling by our methods. The models themselves will be presented elsewhere. Our approach to jet modelling requires fitting to high-fidelity, deep images with linear resolution $\\la$0.25\\,kpc derived from multi-configuration observations at 4.9\\,GHz (C-band) or 8.5\\,GHz (X-band) with the Very Large Array (VLA) at the National Radio Astronomy Observatory\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}. In order to correct the linear polarization for the effects of Faraday rotation (small at these frequencies), we also need to observe at several frequencies in the range 1.3 -- 1.7\\,GHz (L-band). High-quality images of depolarization, rotation measure and spectral index are useful byproducts.  In this paper, we describe: \\begin{enumerate} \\item details of the observations, \\item the source morphologies in total intensity at a range of resolutions, \\item images of the spectral-index distributions and \\item images of the degree of polarization and the apparent magnetic field direction, corrected for Faraday rotation. \\end{enumerate} We also include high-resolution MERLIN\\footnote{MERLIN is a UK National Facility operated by the University of Manchester at Jodrell Bank Observatory on behalf of STFC.} imaging for two of the sources.  Faraday rotation and depolarization are analysed in detail by \\citet[and in preparation]{Guidetti11}.  \\begin{center} \\begin{table} \\caption{The sources. Col.\\,1: name (as used in this paper). Col.\\,2: alternative names. Col.\\,3: redshift. Col.\\,4: linear scale (kpc\\,arcsec$^{-1}$) for our adopted cosmology. Col.\\,5: reference for redshift. Col.\\,6: reference for earlier observations of large-scale radio structure.\\label{tab:sources}} \\begin{minipage}{80mm} \\begin{tabular}{llllll} \\hline Name & Alternative & $z$ & kpc&\\multicolumn{2}{c}{Ref}\\\\ &             &     &arcsec$^{-1}$&O&R\\\\ \\hline NGC\\,193 & PKS\\,0036+03 & 0.0147 & 0.300 &9&4\\\\ &UGC\\,408 &&&&\\\\ 0206+35 & UGC\\,1651 & 0.0377 & 0.748 &3&8\\\\ & 4C\\,35.03 &&&&\\\\ 0755+37 & NGC\\,2484 & 0.0428 & 0.845 &9&1\\\\ M\\,84 & 3C\\,272.1 & 0.0035 & 0.073 &10&5\\\\ &  NGC\\,4374 &&&&\\\\ 3C\\,296 & NGC\\,5532 & 0.0247 & 0.498 &7&6\\\\ 0326+39 &                & 0.0243 & 0.490 &7&2\\\\ \\hline \\end{tabular} References: 1 \\citet{Bondi00}, 2 \\citet{Bridle91}, 3 \\citet{Falco99}, 4 \\citet{Gia11}, 5 \\citet{LB87}, 6 \\citet{LCBH06}, 7 \\citet{Miller02}, 8 \\citet{Morganti87}, 9 \\citet{Ogando}, 10 \\citet{Trager}. \\end{minipage} \\end{table} \\end{center}  We also present an analysis of the spectrum of 3C\\,296 which improves on that given by \\citet{LCBH06} and imaging of the large-scale structure of B2\\,0326+39 to trace its lobe emission closer to the nucleus than in earlier studies. These data complete the documentation of the large-scale structures and spectral-index distributions for all of the lobed FR\\,I sources whose jets we can currently model.  Section \\ref{Obs-red} describes the new observations and data reduction, Section \\ref{Images} presents our results for the sources individually, and Section \\ref{discuss} outlines the phenomenology of their jets and larger-scale emission as a prelude to modelling. Section~\\ref{summary} is a brief summary.  We adopt a concordance cosmology with Hubble constant, $H_0$ = 70\\,$\\rm{km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_\\Lambda = 0.7$ and $\\Omega_M = 0.3$.  \\begin{center} \\begin{table*} \\caption{Journal of VLA observations. Col.\\,1: source name. Col.\\,2: VLA configuration. Col.\\,3: Date of observation. Col.\\,4: centre frequencies for the one or two channels observed (MHz). Col.\\,5 bandwidth (MHz). Col.\\,6: the on-source integration time scaled to an array with all 27 antennas operational. Col.\\,7: VLA proposal code.} \\begin{center} \\begin{tabular}{lclllcl} \\hline Source & Config- &    Date     & $\\nu$ & $\\Delta\\nu$  & t & Proposal \\\\ & uration &             & [MHz]     & [MHz]      & [min] & code\\\\ \\hline NGC\\,193& A & 2007 Jun 02& 4885.1, 4835.1   & 50 &280& AL693 \\\\ & A & 2007 Jun 28& 4885.1, 4835.1   & 50 & 90& AL693 \\\\ & A & 2007 Aug 22& 4885.1, 4835.1   & 50 & 41& AL693 \\\\ & A & 2007 Aug 23& 4885.1, 4835.1   & 50 & 88& AL693 \\\\ & B & 2007 Nov 05& 4885.1, 4835.1   & 50 &318& AL693 \\\\ & B & 2007 Nov 16& 4885.1, 4835.1   & 50 &121& AL693 \\\\ & C & 2008 May 24& 4885.1, 4835.1   & 50 &223& AL693 \\\\ & D & 2007 Mar 11& 4885.1, 4835.1   & 50 & 53& AL693 \\\\ & A & 2007 Jun 28& 1365.0           & 25 & 83& AL693 \\\\ & A & 2007 Aug 22& 1365.0           & 25 & 39& AL693 \\\\ & A & 2007 Aug 23& 1365.0           & 25 & 97& AL693 \\\\ & B & 2007 Nov 16& 1365.0           & 25 &148& AL693 \\\\ & C & 2008 May 24& 1365.0           & 25 & 61& AL693 \\\\ 0206+35 & A & 2008 Oct 13& 4885.1, 4835.1   & 50 &486& AL797 \\\\ & A & 2008 Oct 18& 4885.1, 4835.1   & 50 &401& AL797 \\\\ & B & 2003 Nov 17& 4885.1, 4835.1   & 50 &254& AL604\\\\ & C & 2004 Mar 20& 4885.1, 4835.1   & 50 &88 & AL604\\\\ & A & 2004 Oct 24& 1385.1, 1464.9   & 25 &189& AL604\\\\ & B & 2003 Nov 17& 1385.1, 1464.9   & 25 &110& AL604\\\\ 0755+37 & A &2008 Oct 05& 4885.1, 4835.1   & 50 &477& AL797 \\\\ & A & 2008 Oct 06& 4885.1, 4835.1   & 50 &383& AL797 \\\\ & B & 2003 Nov 15& 4885.1, 4835.1   & 50 & 332& AL604\\\\ & B & 2003 Nov 30& 4885.1, 4835.1   & 50 & 169& AL604\\\\ & C & 2004 Mar 20& 4885.1, 4835.1   & 50 & 125& AL604\\\\ & D & 1992 Aug 2 & 4885.1, 4835.1   & 50 &55  & AM364\\\\ & A & 2004 Oct 25& 1385.1, 1464.9   &12.5&450& AL604\\\\ & B & 2003 Nov 30& 1385.1, 1464.9   &12.5&160& AL604\\\\ & C & 2004 Mar 20& 1385.1, 1464.9   &12.5& 21& AL604\\\\ M\\,84   & A & 1980 Nov 09 & 4885.1           & 50 &223& AL020 \\\\ & A & 1988 Nov 23 & 4885.1, 4835.1   & 50 &405& AW228 \\\\ & A & 2000 Nov 18 & 4885.1, 4835.1   & 50 &565& AW530 \\\\ & B & 1981 Jun 25 & 4885.1           & 50 &156& AL020 \\\\ & C & 1981 Nov 17 & 4885.1           & 50 &286& AL020 \\\\ & C & 2000 Jun 04 & 4885.1, 4835.1   & 50 &138& AW530 \\\\ & A & 1980 Nov 09 & 1413.0           & 25 & 86& AL020 \\\\ & B & 1981 Jun 25 & 1413.0           & 25 & 29& AL020 \\\\ & B & 2000 Feb 09 & 1385.1, 1464.9   & 50 & 30& AR402 \\\\ 0326+39 & D & 1997 Dec 13 & 4885.1, 4835.1 & 50 & 11& AR386 \\\\ & D & 1997 Dec 16 & 4885.1, 4835.1   & 50 & 32& AR386 \\\\ & C & 1998 Dec 4  & 1464.9, 1414.9   & 50 & 11& AR386 \\\\ & D & 1997 Dec 13 & 1464.9, 1385.1   & 50 &  6& AR386 \\\\ & D & 1997 Dec 16 & 1464.9, 1385.1   & 50 &  9& AR386 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:journal} \\end{table*} \\end{center} ",
        "conclusions": "\\label{discuss}  \\subsection{Initial jet propagation} \\label{jets}  There appear to be few, if any, morphological differences between the jet base regions in lobed and tailed FR\\,I sources, whose common properties include the following. \\begin{enumerate} \\item The initial rapid expansion (flaring) and recollimation of the jets is essentially identical in both types of FR\\,I source. \\item Jet bases usually show significant side-to-side asymmetries, although a few very symmetrical examples of each type are known. \\item With the exception of these symmetrical cases, there are further common properties, as follows. \\begin{enumerate} \\item  One jet in each source exhibits a bright region at its base, often with with non-axisymmetric knots and a predominantly longitudinal magnetic field. \\item Jet brightness and polarization asymmetries are correlated: the apparent magnetic field on-axis in the brighter jets is initially longitudinal, but switches to transverse at larger distances; that in the fainter jets is always transverse. \\item The jet/counter-jet ratio falls with increasing distance from the nucleus. \\end{enumerate} \\end{enumerate} These regions at the bases of FR\\,I jets are also those in which our models \\citep{LB02a,CL,CLBC,LCBH06} show that the jets decelerate from relativistic to sub-relativistic velocities. The development of a {\\it transverse} velocity gradient across the jets implies that they decelerate primarily by boundary-layer entrainment of the external medium. We suggest that this entrainment occurs primarily in the dense, kpc-scale coronae of hot plasma that surround the nuclei of twin-jet radio galaxies (probably with an additional contribution from stellar mass loss) and that the entrainment effectively turns off on large scales, as it must to avoid further decollimation. {\\em Chandra} observations have revealed coronae of this type in 0206+35, 0755+37 \\citep{WBH01}, 0326+39 (Hardcastle, private communication), M\\,84 \\citep{M84Chandra} and 3C\\,296 \\citep{Hard05,Croston}; NGC\\,193 may have a similar component \\citep{Gia11}. The coronae for which data are available have central electron densities of $10^5$ to $7 \\times 10^5$\\,m$^{-3}$, central pressures of $3 \\times 10^{-11}$ to $4 \\times 10^{-10}$\\,Nm$^{-2}$ and core radii of 0.3 to 2\\,kpc. It seems likely that lobe plasma is excluded from the immediate vicinity of the nucleus by the high-pressure coronae, so the jets in both types of source initially propagate in essentially identical environments, unshielded from the IGM. Further evolution of the jets will depend on their surroundings: if they propagate through low-density lobe material, then entrainment will effectively cease and they will recollimate to become almost cylindrical flows, as seen in the sources described here. Entrainment rates are also likely to be small in jets without surrounding lobes provided that the external density is low. For example, there is no sign of any lobe surrounding the inner jets of NGC\\,315, yet it has a very small opening angle after recollimation \\citep{CLBC}. This argues for a negligible entrainment rate at distances $\\ga$35\\,kpc, and the external density does indeed fall rapidly on these scales \\citep{Croston}. In contrast, we have argued that entrainment continues at a lower, but still significant rate after recollimation in 3C\\,31, whose inner jets also have no surrounding lobes. In this source, the opening angle after recollimation is larger, our models indicate continuing deceleration and there is a hot-gas component with a large core radius in addition to the inner corona \\citep{LB02b}.  One feature of jet propagation appears to be unique to lobed sources, however. Our high-resolution data for 0206+35 and 0755+37 show that the apparent difference in opening angle between the main and counter-jets seen at lower resolution in these sources is a manifestation of a {\\it two-component} jet structure. In both sources, the main jet and the counter-jet appear to contain both narrow (well-collimated) and broader features on both sides of the nucleus, but the better collimated parts of the main jet are centre-brightened while those in the counter-jet are centre-darkened.  The broader features at the edges of the counter-jets are also slightly brighter than those of the main jets.  What {\\it appears} to be poorer collimation of the counter-jet at low resolution is now seen as a narrow centre-brightened jet opposite a similarly narrow centre-darkened counter-jet, surrounded by broader emission which is slightly brighter around the counter-jet than around the main jet.  We will explore explanations for this ``two-component'' jet structure in terms of relativistic outflow in the well collimated component surrounded by mildly relativistic backflow in the broader component in a later paper (Laing \\& Bridle, in preparation).   \\subsection{Jet termination and lobe structure} \\label{lobes}  With the partial exception of M\\,84, to which we return at the end of this section, the sources described in this paper have lobes similar to those in FR\\,II sources (e.g.\\ \\citealt{AL87,CPDL,K08}).  They exhibit sharp brightness gradients at their outer edges, spectral indices that steepen towards the nucleus from the outer lobes and towards the outer extremities of any lobe, off-axis wings of diffuse emission near their centres, and generally circumferential magnetic fields.  The only obvious difference from the lobes of FR\\,II sources is the evident lack of hot spots at the ends of the FR\\,I jets. Similar spectral gradients occur in a larger sample of lobed FR\\,I sources observed at lower resolution \\citep{Parma99}.  Where the jets dominate the total intensity, they all have similar spectral indices in the range 0.5 to 0.7.  Even when the jets are not obviously dominant in intensity alone, the observed spectral index distributions can trace plausible paths for jets towards the edges of the lobes. This spectral signature implies that steeper-spectrum lobe material has been displaced by flatter-spectrum jet material along an extended pathway through the lobes. Note that the spatial resolution of our data relative to the overall source size is higher than for most published multifrequency imaging of sources in either FR class. Together with the ability to trace the jet flow via its flatter spectrum and the use of intensity gradient images to indicate enhanced compression, our data allow us to identify a number of new types of structures in the lobes.  The jets often terminate in what we have called ``caps'', with the following properties. \\begin{enumerate} \\item They typically occur at the ends of lobes (one example, 0755+37SE, appears recessed, perhaps as a result of projection) and are clearly associated with jet termination. \\item They are bounded at their leading edges by smooth outer isophotes (approximated by segments of circles) with sharp intensity gradients. \\item They also have inner boundaries, again marked by high intensity gradients. \\item Their emission has a flat spectrum, close to that typical of jets ($\\alpha \\approx 0.6$) after accounting for contamination by steeper-spectrum diffuse emission. \\item Four out of five examples are fairly symmetrical with respect to the local jet axis. The exception is 3C\\,296NE. \\end{enumerate}  There are five clear cases of such ``caps'' out of the ten FR\\,I lobes we have studied at high resolution: NGC\\,193E and W, 0206+35NW, 0755+37SE and 3C\\,296NE, and 0755+37NW may well be similar.  Other jet terminations show some, but not all, of the same features. In particular, several show enhanced brightness gradients, but without obvious inner boundaries. In 0206+35SE, the jet bends away from its initial direction and creates a sharp brightness gradient where it impacts on the side of the lobe; the associated emission again has a relatively flat spectrum. Both lobes of 0326+39 probably have similar structures, but the available resolution is not yet high enough to be sure. In 3C\\,296SW and M\\,84N, the jets remain straight until they make oblique impacts on the lobe walls at locations marked by high intensity gradients after which they bend abruptly. The former case also shows flatter-spectrum emission at the impact point. 0755+37NW has a weak ring of emission surrounding a flatter-spectrum region, but this is contained entirely within the outline of the lobe. It is possible that some of these structures (especially the last) are caps seen from an unfavourable angle.  We see little convincing evidence for enhancements in $|\\nabla I|$ {\\em crossing} the jets close to their termination points, such as might be expected from strong shocks. This implies that the flow is internally sub- or transonic. Our estimates of the on-axis flow velocities after the initial rapid deceleration range from $\\la$0.1$c$ to $\\approx$0.6$c$, compared with an internal sound speed of $3^{-1/2}c = 0.58c$ for an ultrarelativistic plasma. Unless there is significant deceleration on scales larger than we model, the implication is that the jets must be very light and energetically dominated by relativistic particles and magnetic field.  The lobes in this class of source are very different from the subsonic, buoyant plumes which are thought to form the outer structures of large FR\\,I sources like 3C\\,31 \\citep{3c31ls}. The picture that emerges for jet termination in lobed FR\\,I sources is that the flow can be traced at least as far as the end of the lobe via its flatter spectrum. Where it impacts on the lobe surface, a high-pressure region (the cap) can be created. The smooth shape of the outer isophotes (compared with the more ragged outline of the lobes) suggests that the forward expansion is at least mildly supersonic with respect to the external medium.  Jet material flows through the cap and back into the lobe, eventually mixing with pre-existing lobe plasma. The flow pattern is sometimes consistent with axisymmetry (or at least appears so in projection) but can bend by large angles without completely losing its collimation.  The clearest example is 3C\\,296NE, where the flow bends by $\\approx$140$^\\circ$ in the plane of the sky and can be traced as far as the trailing edge of the lobe, where its impact is marked by a sharp brightness gradient.  An alternative to the formation of a cap appears to be an oblique collision with the boundary of the lobe. In at least one case, 0206+35SE, the jet deflects from its original straight path before hitting the edge of the lobe. This raises the possibility that the impact point moves around the surface of the lobe, extending it in different directions at different times and lowering the average advance speed of the lobe compared with the instantaneous speed of the impact point, as in the ``dentist's drill'' model of FR\\,II sources \\citep{dentist}.  As noted in Section~\\ref{m84}, M\\,84 appears to be an intermediate case, in which only one of the jets appears to impact on its lobe boundary, but no tails have (yet) developed.  The S lobe of M\\,84 is morphologically very similar to the spurs in 3C\\,31, albeit on a much smaller linear scale. This adds to the developing picture of the transition between fast, well-collimated jets and slow plumes in tailed FR\\,I sources, which is often a two-stage process. The jets first enter bubbles, within which they disrupt, often thrashing around (as in M\\,84S) before disintegrating completely.  The tails are then formed by escape of material from the bubbles along the direction of steepest pressure gradient rather than directly from the jet flow.  A similar morphology is often found in wide-angle tail sources \\citep{HS04}."
    },
    "1107/1107.2396_arXiv.txt": {
        "abstract": "Dust emission is one of the main windows to the physics of galaxies and to star formation as the radiation from young, hot stars is absorbed by the dust and reemitted at longer wavelengths. The recently launched Herschel satellite now provides a view of dust emission in the far--infrared at an unequaled resolution and quality up to 500~$\\mu$m. In the context of the Herschel HERM33ES open time key project, we are studying the moderately inclined Scd local group galaxy M33 which is located only 840~kpc away. In this article, using Spitzer and Herschel data ranging from 3.6~$\\mu$m to 500~$\\mu$m, along with HI, H$\\alpha$ maps, and GALEX ultraviolet data we have studied the emission of the dust at the high spatial resolution of 150~pc. Combining Spitzer and Herschel bands, we have provided new, inclination corrected, resolved estimators of the total infrared brightness and of the star formation rate from any combination of these bands. The study of the colors of the warm and cold dust populations shows that the temperature of the former is, at high brightness, dictated by young massive stars but, at lower brightness, heating is taken over by the evolved populations. Conversely, the temperature of the cold dust is tightly driven by the evolved stellar populations. ",
        "introduction": "The infrared (IR) is one of the main windows to the physics of galaxies. Not only does it give access to the stellar content of galaxies in the near--IR \\citep{bell2001a}, it also provides us with insight into the fundamental processes of star formation. Indeed, as dust absorbs the energetic radiation from massive young stars, it reemits this energy in the IR. In other words, the IR traces star--formation.  The advent of IR telescopes such as IRAS (Infrared Astronomical Satellite) and its successors ISO (Infrared Space Observatory), Spitzer, Akari, and now WISE (Wide-field Infrared Survey Explorer) and Herschel has revolutionized our view on the universe with the discovery of a large number of dust-enshrouded galaxies whose existence had never been suspected. The mid-- and far--IR have now become a cornerstone to our understanding of galaxy formation and evolution, accounting for a third of the bolometric luminosity in quiescent galaxies such as the Milky Way, 70\\% in more actively star--forming galaxies such as M82, and up to 99\\% for ULIRGs \\citep[Ultra Luminous Infrared Galaxies,][]{lagache2005a}.  The recently--launched Herschel Space Observatory provides a continuous coverage from 70~$\\mu$m to 500~$\\mu$m at an unprecedented resolution and sensitivity for the longer wavelengths. It probes dust--enshrouded star--formation across cosmic time, from the local group to high--redshift galaxies. Its exquisite resolution (from $\\sim6$\" at 70~$\\mu$m to $\\sim 38$\" at 500~$\\mu$m) makes it possible to resolve and study individual star--forming regions within nearby galaxies \\citep[][for instance]{beirao2010a,bendo2010a,boquien2010b,cormier2010a,verley2010b}. For galaxies at cosmological distances, deep IR surveys are key tools for understanding the physics of galaxy formation and evolution \\citep[][and many others]{chary2001a,dale2001a,dale2002a}. The sensitivity of Herschel allows us to probe star--formation in the IR down to the regime of quiescent star--forming galaxies at high redshift as in the case of the Herschel--GOODS deep survey. It is therefore not only timely but also of critical importance to understand the origin of the IR emission in such galaxies in the nearby universe to interpret properly these observations.  Dust emission in galaxies can have several unrelated origins as for instance suggested by \\cite{sauvage1994a,kewley2002a}. (1) Massive stars in star--forming regions heat and excite the dust, as described above. This emission is generally compact and corresponds to clumps of star--formation within galaxies. (2) A large scale diffuse emission may arise due to cirrus \\citep{helou1986a} being heated by (a) energetic radiation escaping from individual star--forming regions \\citep{sauvage1990a,xu1990a,popescu2000a,misiriotis2001a,popescu2002a,popescu2005a,tabatabaei2007a} and/or (b) the general radiation field of evolved stars \\citep{helou1986a,xu1996a,li2002a,boselli2004a}. (3) Hot grains in the photosphere or circumstellar atmosphere of mass--losing stars also contribute, especially in the mid--IR \\citep[][for instance]{jura1987a,knapp1992a,mazzei1994a}. (4) Embedded AGN affect the surrounding dust \\citep{degrijp1985a,wu2007a}.  Observations have shown that the dust forms as soon as metals are available leading to some absorption by the dust at relatively high redshift \\citep{meurer1997a,giavalisco2004a,bouwens2009a}. This means that dust also reprocesses the radiation from young stars at cosmological distances and, as a consequence, IR emission acts as a star--formation tracer even at high redshift. This hints at the fundamental importance of the IR not only to study star--formation in nearby galaxies but also in the distant universe. For a decade, studies have used the IR luminosity function to explore the formation and evolution of galaxies, providing an extinction--free estimate by essence of the cosmic star formation rate density. It serves as a test bench for models of structure formation and evolution \\citep[e.g.,][]{kay2002a}.  The use of the IR as a SFR (star formation rate) estimator is based on three assumptions. First, it is assumed that dust opacity is infinite throughout the galaxy. That is, all the energetic radiation from young stellar populations is absorbed and reemitted in the IR. While this assumption is fulfilled in LIRGs and ULIRGs, in more quiescent galaxies it is not the case. The combination of the IR continuum luminosity with an optical recombination line or the UV (ultraviolet) luminosity permits to circumvent this problem \\citep{calzetti2007a,kennicutt2007a,leroy2008a,bigiel2008a,kennicutt2009a}. The second assumption is that the young stellar populations dominate the radiation field in extinction--sensitive bands (UV and optical). In other words it means that AGN heating as well as heating from evolved stellar populations and dust emission in circumstellar envelopes are small compared to it. This question is of crucial importance as it determines whether a given IR band is a reliable star--formation tracer or not. As the diffuse emission can represent up to nearly 90\\% of the IR flux in Sa galaxies \\citep{sauvage1992a}, it can lead to an error on the SFR up to an order of magnitude depending on its origin. Finally the last assumption is that the initial mass function is universal.  There is an active and on--going debate in the literature regarding the actual process giving rise to the diffuse emission. Its presence has been unambiguously detected in numerous studies since \\cite{lonsdale1987a}. The non star--forming related contamination is non--negligible for whole galaxies according to some authors, however its amplitude is not assessed precisely, with some studies on individual galaxies and samples of galaxies finding conversely that star--formation drives the majority of the IR luminosity \\citep{sauvage1992a,devereux1994a,devereux1994b,devereux1995a,buat1996a,israel1996a,walterbos1996a,devereux1997a,jones2002a,kewley2002a,bell2003a,boselli2004a,forster2004a,hinz2004a,cannon2006b,perez2006a,calzetti2007a,montalto2009a,dacunha2010a,calzetti2010a}. These results hint at a variation from galaxy to galaxy in a more fundamental way than the relation with galaxy type found by \\cite{sauvage1992a}. When studying individual star--forming regions, it appears that the warm dust peaking around 70~$\\mu$m is a good tracer of star formation as was found both in the Magellanic Clouds \\citep{lawton2010a} and in star forming regions in nearby galaxies \\citep{li2010a}.  Here we use Herschel observations taken in the context of the HERM33ES open time key project \\citep{kramer2010a} in combination with Spitzer IRAC (Infrared Array Camera) and MIPS (Multiband Imaging Photometer for Spitzer) data as well as GALEX FUV (far--ultraviolet), and ground--based H$\\alpha$ maps of a local group galaxy, in order to characterize the emission from the warm and cold dust populations\\footnote{The case of the PAH properties will be treated extensively in Rosolowsky et al. (in preparation). We concentrate here primarily on the emission of the warm and cold dust components.}. By convention, we designate as ``warm dust'' the population that peaks typically at $\\mathrm{\\sim50~K}$ and by ``cold dust'' the one that peaks at $\\mathrm{\\sim20~K}$. The selected target is M33, a nearby Scd galaxy located only 840~kpc away \\citep{freedman1991a} with an inclination of 56$^{\\circ}$ \\citep{regan1994a}. It provides us with an exceptional physical resolution better than $\\sim150$~pc at 500~$\\mu$m allowing us to resolve the various morphological components of the galaxy. Finally, the absence of an active nucleus excludes one possible dust heating mechanism and a shallow metallicity gradient limits the influence of the radial evolution of the metallicity on the emission of the dust.  In Sec.~\\ref{sec:obs} we present the observations and the  data reduction. Sec.~\\ref{sec:results} provides the results, which are discussed in Sec.~\\ref{sec:discussion}. The main conclusions are given in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Using Spitzer and Herschel IR data from 3.6~$\\mu$m to 500~$\\mu$m, along with GALEX FUV data and HI and H$\\alpha$ maps, we have studied the properties and the origin of the warm and cold dust emission in the moderately inclined, local group galaxy M33, at a spatial resolution of 150~pc. To do so, we have studied the evolution of the dust colors both as the position in the galaxy (including the radial distance) and the brightness. We have found the following results:  \\begin{enumerate} \\item The colors of the warm (as seen in the 24~$\\mu$m to 160~$\\mu$m bands) and cold dust (observed from 250~$\\mu$m to 500~$\\mu$m) components seem to be predominantly driven by the radiation field intensity. \\item Combining any set of Spitzer and Herschel bands, we have provided correlations to estimate both the TIR brightness and the SFR, extending the results of \\cite{boquien2010a,boquien2010b,verley2010b}, which is of importance for the study of high redshift galaxies. \\item The color trends of the warm and the cold dust show that they are heated by different sources. At higher SFR, the warm dust temperature seems to be driven by star formation. As star formation decreases, the temperature is increasingly driven by another component, most likely the evolved stellar populations. The cold dust temperature seems to be driven by the old stellar population, with a tight correlation with the local stellar mass. \\end{enumerate}  This paper lays the basis of additional studies in the context of the HERM33ES project which will go in much more detail on some of the points elaborated here. Upcoming papers will provide further insight on the PAHs emission (Rosolowsky et al., in preparation), the local dust temperature and mass (Xilouris et al., in preparation), the dust properties (Tabatabaei et al., in preparation), the HII regions (Rela\\~no et al., in preparation), the local physical conditions using PACS spectroscopy (Mookerjea et al., 2011, submitted), etc."
    },
    "1107/1107.5325_arXiv.txt": {
        "abstract": "{Searching for extrasolar planets through radial velocity measurements relies on the stability of stellar photospheres. Several phenomena are known to affect line profiles in solar-type stars, among which stellar oscillations, granulation and magnetic activity through spots, plages and activity cycles.} {We aim at characterizing the statistical properties of magnetic activity cycles, and studying their impact on spectroscopic measurements such as radial velocities, line bisectors and line shapes.} {We use data from the HARPS high-precision planet-search sample comprising 304 FGK stars followed over about 7 years. We obtain high-precision Ca II H\\&K chromospheric activity measurements and convert them to $R'_{\\mathrm{HK}}$ indices using an updated calibration taking into account stellar metallicity. We study $R'_{\\mathrm{HK}}$ variability as a function of time and search for possible correlations with radial velocities and line shape parameters.} {The obtained long-term precision of $\\sim$0.35\\% on S-index measurements is about 3 times better than the canonical Mt Wilson survey, which opens new possibilities to characterize stellar activity. We classify stars according to the magnitude and timescale of the Ca II H\\&K variability, and identify activity cycles whenever possible. We find that 39$\\pm$8\\% of old solar-type stars in the solar neighborhood do not show any activity cycles (or only very weak ones), while 61$\\pm$8\\% do have one. Non-cycling stars are almost only found among G dwarfs and at mean activity levels $\\log{R'_{\\mathrm{HK}}}$ $<$ -4.95. Magnetic cycle amplitude generally decreases with decreasing activity level. A significant fraction of stars exhibit small variations in radial velocities and line shape parameters that are correlated with activity cycles. The sensitivity of radial velocities to magnetic cycles increases towards hotter stars, while late K dwarfs are almost insensitive.} {Activity cycles do induce long-period, low-amplitude radial velocity variations, at levels up to $\\sim$25 m\\,s$^{-1}$. Caution is therefore mandatory when searching for long-period exoplanets. However, these effects can be corrected to high precision by detrending the radial velocity data using simultaneous measurements of Ca II H\\&K flux and line shape parameters.} ",
        "introduction": "The advent of large-scale Doppler surveys to search for extrasolar planets around FGKM stars in the solar neighborhood has produced an impressive body of high-resolution spectroscopic data extending over the past 20 years or so. Precise radial velocities are obviously the main products that are derived from these data, but more generally the existing libraries of spectra make it possible to study in details many properties of solar-type stars. The knowledge of extrasolar planets has always been intimately related to the knowledge of their parent stars. An important example is the need for precise fundamental stellar parameters such as effective temperature, metallicity, mass and radius to properly derive exoplanet parameters. The behavior of stellar photospheres also plays a crucial role: since the discovery of 51 Peg b \\citep{mayor95}, questions have arisen about the effects of stellar photospheric \"jitter\" on the detection and characterization of exoplanets \\citep[e.g.][]{Baliunas97}. The outer convective envelope that is present in solar-type stars gives rise to several phenomena that can have an impact on derived exoplanet properties, through variations in measured disk-averaged radial velocity and photometric flux. In order of increasing timescales, one can mention p-mode oscillations \\citep{christensen04}, granulation and supergranulation \\citep{harvey85}, magnetic activity inducing surface inhomogeneities rotating with the star \\citep{saar97}, and finally magnetic activity cycles over timescales of years and decades \\citep{baliunas95}. From a stellar physics point of view, all these phenomena are obviously interesting in their own right since they can reveal the details of stellar internal structure and dynamo mechanisms at play in solar-type stars. Moreover, the statistical study of stellar activity in a large number solar-type stars can also inform us about the past and future evolution of our Sun, and its potential impact on Earth climate.  With the ever-increasing precision of radial velocity (RV) and transit observations of exoplanets, the need to better understand the host stars has become more acute in recent years. In particular, radial velocity surveys are now probing the sub-m\\,s$^{-1}$ regime \\citep{pepe11}, a level at which stellar activity perturbations become non-negligible even for quiet stars and must be monitored. The study of the effects of star spots on radial velocities has a long history, see e.g. \\citet{saar97,saar98,santos00,saar00,queloz01}. More recently, the field has received closer attention again following the discovery of transiting planets around relatively active stars like CoRoT-7 \\citep{queloz09,boisse11,hatzes11} or M dwarfs like GJ~436 \\citep{knutson11}. At the same time, efforts to characterize RV jitter in quiet solar-type stars, in particular the Sun itself, have also intensified. For example, \\citet{dumusque11a,dumusque11b} simulate the impact of p-mode oscillations, granulation, supergranulation and active regions on RV data, based on asteroseismic observations of a few solar-type stars. \\citet{lagrange10} and \\citet{meunier10} use solar irradiance and Ca II archival data to infer the level of RV jitter caused by spots and plages. While these studies offer important insights into the causes of RV jitter, they rely at least partly on unverified assumptions to relate a given physical phenomenon at the stellar surface to the corresponding RV perturbation, as actually measured by RV planet-search techniques. The ultimate \"truth\" in this domain must necessarily come from actual measurements of precise RVs, photometry and/or spectroscopic diagnostics in the form of densely-sampled time series covering the relevant timescales. The present paper is an attempt is this direction, focusing on the longer timescales.  Over years and decades, the main cause of variability in the Sun is the occurrence of a magnetic cycle with a quasi-periodicity of 11 years (or 22 years depending on definition). The study of this cycle has a long history, starting with its discovery by S. H. Schwabe in 1843. The question whether a similar phenomenon also occurs in other stars was addressed in the 1960s by O. Wilson, who started what would become the famous Mt Wilson program for measuring chromospheric emission in stars of the solar neighborhood \\citep{wilson68,wilson78}. The Mt Wilson spectrophotometers regularly monitored the flux in the core of the Ca II H \\& K lines in the sample stars for more than 30 years, yielding what is still today the reference database for magnetic activity in solar-type stars. Long-term results from this program were published in several papers \\citep[e.g.][]{wilson78,duncan91,baliunas95}, which demonstrated that activity cycles similar to the solar one are indeed common in other stars.  In the meantime, other stellar activity surveys, and the development of large-scale Doppler exoplanet programs, have also produced a large quantity of Ca II H \\& K observations, and several studies have been conducted to investigate magnetic activity levels in solar-type stars \\citep[e.g.][]{henry96,santos00,tinney02,wright04a,hall07,isaacson10,santos10}. All these \"second-generation\" projects tried to reproduce as closely as possible the method used by the Mt Wilson survey to measure Ca II H \\& K emission, i.e. by computing the so-called S index which normalizes the measured Ca II core flux by the flux in two \"continuum\" bandpasses on the blue and red sides of the Ca II lines. The Mt Wilson scale, as defined by the S index, has therefore become the standard way of measuring chromospheric activity in stars.  Most projects have actually focused on the mean activity level of stars, and not on the time variability of the activity. To the best of our knowledge, only \\citet{baliunas95}, using the Mt Wilson data, have actually performed a systematic search for magnetic cycles and derived their main parameters such as cycle period. In the present paper we attempt to do a similar study to better characterize overall cycle properties in older solar-type stars, and compare them to the Mt Wilson survey. Then we search for the effects of magnetic cycles on the shapes of photospheric lines, with the ultimate goal of disentangling stellar RV \"jitter\" from real barycentric motions of the star.  The idea that magnetic cycles may influence spectral lines dates back at least to \\citet{dravins85}. The rationale behind this is that the convective patterns at the surface of solar-type stars may change along the magnetic cycle under the influence of changing magnetic field strength. Indeed, the convection is greatly reduced in active regions, which causes a decrease in the convective blueshift usually exhibited by photospheric lines \\citep[see][and references therein for more detailed studies of these effects]{dravins82,livingston82,brandt90,gray92,lindegren03,meunier10}.  A few attempts to measure such an effect in the Sun yielded somewhat contradictory results: indeed, \\citet{deming94} find a peak-to-peak RV variation of 28 m\\,s$^{-1}$ over the solar cycle, while \\citet{mcmillan93} obtain constant RVs within $\\sim$4 m\\,s$^{-1}$. We note here that the two studies used very different spectral lines for their measurements (CO lines in the IR vs. atomic UV lines). Measurements of similar RV effects in other stars have remained inconclusive to date \\citep[e.g.][]{campbell88,saar00, santos10}, although \\citet{santos10} find clear correlations in several stars between magnetic cycles and line shape parameters like bisector and line width.  In the present paper we use high-precision data from the HARPS spectrograph \\citep{mayor03}, obtained in the context of a planet-search survey around FGK dwarfs in the solar neighborhood. This survey focuses on the search for low-mass planets, i.e. Neptunes and super-Earths, which induce RV variations of only a few m\\,s$^{-1}$ on their parent star \\citep[see e.g.][]{lovis06a,mayor09b,lovis11,pepe11}. Potential effects of magnetic cycles at this level could induce spurious long-term drifts in the data that may perturb the detection of low-mass objects at shorter periods due to limited sampling, or even mimic the signal of giant planets orbiting at several AUs from the star. It is therefore crucial to understand the effects magnetic cycles may have on our sample stars. In an accompanying paper to the present study, \\citet{dumusque11c} show several examples of planetary systems around stars showing magnetic cycles, and discuss how to correct the perturbing effects of these. ",
        "conclusions": "\\subsection{Calibration of $R'_{\\mathrm{HK}}$ taking into account metallicity}  In this paper we provide an updated calibration of the $R'_{\\mathrm{HK}}$ index, which takes into account the effects of metallicity on stellar bolometric luminosity. Based on a calibration of effective temperature as a function of $B-V$ color and metallicity, a correction factor to the usual $R'_{\\mathrm{HK}}$ definition can be applied and removes the observed trend of increasing $R'_{\\mathrm{HK}}$ values with decreasing metallicity. The correction yields a sharper distribution of mean $R'_{\\mathrm{HK}}$ values for our sample stars, confirming that stars of different metallicities were exhibiting an artificial spread in $R'_{\\mathrm{HK}}$ with the uncorrected calibration.   \\subsection{Statistics of magnetic cycles}  The present survey has allowed us to revisit some of the properties of magnetic cycles in solar-type stars. To the best of our knowledge, most of the information previously available on this subject comes from the Mt Wilson project, and in particular the global analysis of \\citet{baliunas95} \\citep[see however][for a recent study]{saar09}. Two main difficulties arise when aiming at characterizing magnetic cycles: the required long-term precision on Ca II H\\&K measurements, and the necessity to monitor many stars over decadal timescales. It turns out that large radial-velocity planet-search surveys, which started in the 1990s, can actually provide both, since they are based on high-resolution echelle spectrographs and regularly monitor hundreds of solar-type stars over many years. Several surveys have already yielded mean activity levels for many stars, but {\\em time series} of activity measurements have been rarely discussed.  The HARPS data presented here are of excellent quality in terms of precision on $R'_{\\mathrm{HK}}$, but they span \"only\" 7-8 years, compared to 25 years for the published Mt Wilson data \\citep{baliunas95}. However, as shown in Sect.~\\ref{SectMagneticCycles}, this did not prevent us from detecting magnetic cycles with estimated periods up to $\\sim$15 yr, and long-term drifts of undetermined periodicity. Obviously, the precision on fitted periods quickly decreases when going beyond the time span of observations, but even uncertainties of 25-50\\% on the period still provide useful information on cycle properties.  An important simplification adopted in this work is the assumption of perfectly sinusoidal magnetic cycles \\citep[as in][]{baliunas95}. We know from the Sun that they are in fact not really sinusoidal for any given cycle number, and moreover their amplitudes and shapes may vary significantly from one cycle to the next. However, our survey usually covers only one or two cycles, so that we are not able to study the {\\em evolution} of the magnetic cycle in any given star. We rather obtain a random snapshot of cycles in many stars at the same time. With these limitations in mind, we can summarize the main results of this work as follows:  \\begin{itemize}  \\item The global occurrence of magnetic cycles with semi-amplitude $A$ $>$ 0.04 is 61$\\pm$8\\% among older solar-type stars. Correspondingly, 39$\\pm$8\\% of older solar-type stars show a flat activity record. These numbers are roughly compatible with the results of \\citet{baliunas95}, although these authors do not explicitly provide comparable estimates and use a different classification scheme. The picture is markedly different for subgiants, for which the occurrence of magnetic cycles drops to only 12$\\pm$5\\%.  \\item The distribution of magnetic cycle periods shows a broad maximum between 2-10 yr, and a sharp decrease towards longer periods. This may be due in part to the limited duration of the survey, but the number of observed long-term trends does not seem to be sufficient to produce a large number of long-period cycles. The decrease towards longer periods therefore seems to be real. There may be some discrepancies between our period distribution and the one of \\citet{baliunas95}, in the sense that we detect more short-period cycles than they do. It must be noted, however, that short-period cycles are often less well-defined (more dispersion and evolution with time) than longer-period ones, and the discrepancy may be explained by a different classification scheme combined with a longer duration of observations (\\citet{baliunas95} classify as \"variable\" the stars with poorly defined periodicities). On the other hand, both studies seem to agree on the low number of cycles with $P$ $\\gtrsim$ 15 yr. This confirms that the limited duration of our survey does not significantly affect the derived overall properties of magnetic cycles.  \\item The distribution of magnetic cycle semi-amplitudes shows a broad peak between zero and 0.2 in $R'_{\\mathrm{HK}}$. Given the observational bias against low amplitudes, it seems that magnetic cycles significantly less pronounced than the solar one are the rule rather than the exception.  \\item We find no clear correlations between cycle periods and amplitudes, nor between cycle periods and mean activity levels. The sometimes large error bars on cycle periods prevent us from going into further details. The occurrence and stability in time of short-period cycles ($P$ $\\lesssim$ 5 yr) should be further investigated, together with methods to disentangle them from evolution of active regions on rotational timescales.  \\item We find a clear trend of increasing cycle amplitudes with increasing mean activity levels. In particular, no stars with mean $\\log{R'_{\\mathrm{HK}}}$ $\\approx$ -5.0 show cycles as pronounced as the Sun. We also find that the upper envelope of cycle amplitudes shows a peak in the temperature range 5000-5600\\,K, while hotter stars do not exhibit very large amplitudes.  \\item The distribution of mean $\\log{R'_{\\mathrm{HK}}}$ values is markedly different between G and K dwarfs, in the sense that G dwarfs tend to accumulate around -5.0, while K dwarfs are more spread between -4.7 and -5.0. A possible explanation is the slower decrease in activity with age in K dwarfs, which need more than $\\sim$6-8 Gyr to reach a mean level of -5.0.  \\item Non-cycling stars are only found at mean activity levels lower than about -4.95, and almost only among G dwarfs. Cycling stars are preferentially found at higher activity levels, although some of them are also found around -5.0. The number of well-observed K dwarfs below -4.95 is too low to properly study the occurrence of non-cycling stars in this region of parameter space.  \\item A possible interpretation of these results is that magnetic cycle amplitudes decrease with age, and tend to disappear when stars reach a mean activity level of about -5.0. This is also supported by the quasi-absence of cycles in subgiants.  \\item The Sun, with a mean $\\log{R'_{\\mathrm{HK}}}$ of -4.91, is still in the activity range where most stars exhibit magnetic cycles. One can conjecture that solar cycles will slowly disappear as the mean activity decreases to -5.0, possibly within a few Gyr. It is difficult to estimate from our data how frequent periods of extended minima like the Maunder minimum are for any given star, since we only have instantaneous snapshots of the state of many stars. What can be said is that the activity level of the Sun was almost certainly below -4.95 during the Maunder minimum since we do not find any non-cycling stars above that threshold. Considering the well-established decrease in activity with age in solar-type stars, it is more likely that the \"clump\" of G dwarfs around -5.0 is mainly the signature of an age effect, and not a population of stars in temporary Maunder-minimum states. However, an unknown proportion of these stars are probably in such a state. This also supports the view that Maunder-minimum stars cannot be found on the basis of their mean activity level only \\citep{wright04b}, and should rather be identified by the demonstrated absence of a magnetic cycle.  \\end{itemize}   \\subsection{Effects of magnetic cycles on spectral lines and radial velocities}  It has long been hypothesized that magnetic cycles have an impact on photospheric spectral lines \\citep[e.g.][]{dravins85,gray92,lindegren03}. Here we provide a global assessment of magnetically-induced variability in four quantities derived from the HARPS cross-correlation function (RV, FWHM, contrast and BIS), which is constructed as a \"master\" spectral line built from thousands of mainly weak to moderate individual Fe lines.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{plot_effective_RV_impact.eps} \\caption{Effective impact of magnetic cycles on RV for all sample stars with a detected cycle, shown as a function of $T_{\\mathrm{eff}}$.} \\label{FigPlotEffectiveRVImpact} \\end{figure}  We observe long-term correlations between $R'_{\\mathrm{HK}}$, RV, FWHM, contrast and BIS in many stars of our sample, with varying sensitivities between stars of different spectral types and various behaviors for different quantities. Globally speaking, for a given variation in $R'_{\\mathrm{HK}}$, we find that RV and BIS are more sensitive in G dwarfs than in K dwarfs, whereas the opposite is true for the FWHM. The contrast sensitivity is roughly constant with spectral type, but varies significantly with metallicity. Overall, sensitivity variations with temperature are more important than with metallicity for RV, FWHM and BIS. The sensitivity of RV to magnetic cycles is almost zero at $T_{\\mathrm{eff}}$ $\\approx$ 4800 K, while the same is true for FWHM at $T_{\\mathrm{eff}}$ $\\approx$ 6200 K. However, the effective impact of a magnetic cycle for any given star is obtained by multiplying these sensitivities with the actual $R'_{\\mathrm{HK}}$ semi-amplitude, which may be large and therefore have a non-negligible effect even if the sensitivity is low. We show in Fig.~\\ref{FigPlotEffectiveRVImpact} the actual impact of magnetic cycles on the RV data for all stars in the sample with a detected cycle. For each star the RV sensitivity is computed from the model given in Eq.~\\ref{EqSensitivities}, and the result is multiplied by the measured $R'_{\\mathrm{HK}}$ semi-amplitude $A$. We see that induced RV semi-amplitudes can reach up to 11 m\\,s$^{-1}$ in the worst case. More typical values are between 0 and 3 m\\,s$^{-1}$ (there are unaffected stars at all spectral types). Interestingly, the model predicts an {\\em anti-correlation} between $R'_{\\mathrm{HK}}$ and RV (i.e. a negative RV semi-amplitude in Fig.~\\ref{FigPlotEffectiveRVImpact}) for stars cooler than $\\sim$4800 K. At least in one case this surprising result seems to be confirmed by an in-depth analysis: the star HD\\,85512 \\citep{pepe11}. Looking at Fig.~\\ref{FigPlotEffectiveRVImpact}, we also see that despite the drop in sensitivity towards cool stars, one can still have RV effects of several m\\,s$^{-1}$ because of the occurrence of strong magnetic cycles in K dwarfs.  The previous study that is most directly comparable to the results presented here is the one by \\citet{santos10}. These authors obtained long-term HARPS data on eight stars known to have clear magnetic cycles, and studied the behavior of CCF parameters in the same way as we do here. Clear correlations were found between the S index, FWHM and BIS data, while the contrast was anti-correlated to the S index. These findings are in excellent agreement with the present results. On the other hand, attempts to find correlations with the RV data remained inconclusive, or actually RV effects were too small to be detected. The sensitivity model derived in this paper provides a natural explanation for this lack of correlations: it turns out that the five cycling stars actually studied by \\citet{santos10} span a very narrow range of effective temperatures (200 K), centered at 5000 K. In this regime RV sensitivities are close to zero, and the S index variations were not large enough to induce a measurable RV effect. The only hot star in the sample, HD\\,216385, is a magnetically-constant star that was chosen as a standard. Two other, higher-temperature stars in the sample could not be studied because of strong short-term variability and binarity. We conclude that it is only bad luck that prevented \\citet{santos10} from finding correlations between RV and S index.  In general, the RV sensitivities to magnetic cycles are such that the induced RV effects can be of the same order of magnitude and period as the signal of Jupiter-like planets orbiting at several AUs from the star. Therefore, it is critical to check the behavior of the chromospheric activity index before announcing any long-period planet with a RV semi-amplitude lower than $\\sim$15 m\\,s$^{-1}$. It would be extremely useful for all planet-search teams to measure chromospheric activity (even at lower precision) in all stars exhibiting such low-amplitude, long-period RV signals, and to check whether any correlations are found with the magnetic cycle. This obviously also applies to already published planets, e.g. the Jupiter-analog candidate around HD\\,154345 \\citep{wright08}. If simultaneous activity measurements are available, then the present work shows that they can be easily used to correct RV data from magnetic cycle effects. Two approaches can be used: either one includes an additional sinusoidal component in the Keplerian model, fixing its period and phase to the ones of the magnetic cycle, or one can use the sensitivity model presented here to {\\em a priori} correct the RV data based on a fit of the magnetic cycle on the activity data. Several examples of these correction procedures and their accuracies can be found in \\citet{dumusque11c}.  The sensitivity model also allows us to estimate the RV effect that is expected in the Sun due to the solar magnetic cycle. With $T_{\\mathrm{eff}}$ = 5770 K, [Fe/H] = 0 and $A$ = 0.26, we obtain a RV semi-amplitude of 8.2 m\\,s$^{-1}$ and thus a peak-to-peak effect of about 16 m\\,s$^{-1}$. This is in qualitative agreement with both theoretical estimates \\citep[e.g.][]{lindegren03} and a previous attempt to actually measure it \\citep{deming94}. We note that the effect is likely to depend on the particular set of spectral lines being considered, so that the comparison with literature values is not straightforward. This could potentially explain the null result of \\citet{mcmillan93}. A more recent attempt to predict this same effect and its impact on exoplanet searches was published by \\citet{meunier10}. They obtain a full amplitude over the solar cycle of about 8 m\\,s$^{-1}$, which is lower than our result but can also be considered as qualitatively in agreement given the differences in the details between our observations and their simulations. We note in passing that \\citet{meunier10} went on to simulate planet detection limits in the presence of stellar noise, and concluded that the effects of magnetic cycles would make the detection of Earth-like planets impossible. We emphasize here that these long-term effects are actually rather easy to model and correct using simultaneous $R'_{\\mathrm{HK}}$ measurements. Their planet detection limits are therefore likely far too pessimistic.  Finally, we briefly discuss theoretical explanations for the observed effects of magnetic cycles on spectral lines. As already argued by several authors, the main phenomenon at play here is likely to be the \"freezing\" effect of magnetic fields on granular motions in the outer convective zone of solar-type stars \\citep[e.g.][]{livingston82,brandt90}. The stronger the magnetic field, the slower the convective motions. As a consequence, the usual convective blueshift of spectral lines, caused by the temperature-velocity correlation in granules \\citep{dravins82,kaisig82}, is decreased in active regions compared to the quiet photosphere. Therefore, stellar lines will appear redshifted when the average (or local) magnetic field strength is higher. This may explain the observed RV - $R'_{\\mathrm{HK}}$ positive correlation over the magnetic cycle. However, other effects may also play a role, such as changes in large-scale surface flows over the magnetic cycle \\citep{makarov10}, or perhaps changes in global mean effective temperature as the total spot/plage coverage varies (sensitivity of spectral lines to temperature).  Differential effects as a function of effective temperature or metallicity are more difficult to explain. Convective motions are slower in cooler stars, diminishing the amount of convective blueshift and therefore also reducing possible variations with magnetic field strength. The strong increase in FWHM sensitivity towards cooler temperatures should also be considered to help us understand the effect. The direct impact of magnetic fields on spectral lines, through Zeeman broadening, should be investigated. 3D hydrodynamical simulations of stellar atmospheres including magnetic fields are probably necessary to better understand all these observational findings. On the observational side, the lists of spectral lines that are used to build the cross-correlation function for precise RV measurements should be carefully checked and optimized to either minimize undesirable RV effects, or actually maximize them in order to better understand the underlying physics affecting line shapes over magnetic cycles.  In conclusion, the present study has allowed us to make important progress in the knowledge of magnetic cycles in solar-type stars in general, and in their effects on spectral lines in particular. Not only are these results hopefully useful in their own right for stellar physics, but they are also of great help for the search for exoplanets with the RV method. Understanding stellar RV \"jitter\" is the single most important challenge in the quest for Earth-like planets in the habitable zones of solar-type stars. The excellent quality of the HARPS data, and the different diagnostics available from the spectra, make it possible to monitor stars with unprecedented scrutiny. Further studies of the impact of magnetic activity at all timescales are ongoing and will allow us to disentangle even better stellar activity and planetary signals."
    },
    "1107/1107.0736_arXiv.txt": {
        "abstract": "We investigate the dynamical evolution of hierarchical three-body systems under the effect of tides, when the ratio of the orbital semi-major axes is small and the mutual inclination is relatively large (greater than $20^\\circ$). Using the quadrupolar non-restricted approximation for the gravitational interactions and the viscous linear model for tides, we derive the averaged equations of motion in a  vectorial formalism  which is suitable to model the long-term evolution of a large variety of exoplanetary systems in very eccentric and inclined orbits. In particular, it can be used to derive constraints for stellar spin-orbit misalignment, capture in Cassini states, tidal-Kozai migration, or damping of the mutual inclination. Because our model is valid for the non-restricted problem, it can be used to study systems of identical mass or for the outer restricted problem, such as the evolution of a planet around a binary of stars. Here, we apply our model to three distinct situations: 1) the HD\\,80606 planetary system, for which we obtain the probability density function distribution for the misalignment angle, with two pronounced peaks of higher probability around 53$^\\circ$ and 109$^\\circ$; 2) the HD\\,98800 binary system, for which we show that initial prograde orbits inside the observed disc may become retrograde and vice-versa, only because of tidal migration within the binary stars; 3) the HD\\,11964 planetary system, for which we show that tidal dissipation combined with gravitational perturbations may lead to a decrease in the mutual inclination, and a fast circularization of the inner orbit. ",
        "introduction": "At present, about 50 multi-planet systems have been reported, out of which roughly 1/3 possess close-in planets with semi-major axis smaller than 0.1\\,AU. There are also indications that about half of these stars are likely to have distant companions \\citep{Fischer_etal_2001}. {\\bfx In addition, close binary star systems (separation smaller than 0.1\\,AU) are often accompanied by a third star \\citep[e.g.][]{DAngelo_etal_2006}. } Therefore, although hierarchical systems are considerably different from our Solar system, they represent a significant fraction of the already known systems of stars and planets. These  systems are particularly interesting from a dynamical point of view, as they can be stable for very eccentric and inclined orbits and thus present uncommon behaviors. In particular, they become interesting when the two innermost bodies are sufficiently close to undergo significant tidal interactions over the age of the system, since the final outcome of the evolution can be in a configuration that is totally different from the initial one.  The origin and evolution of the orbital configurations of multi-body systems can be analyzed with direct numerical integrations of the full equations of motion, but the understanding of the dynamics often benefits from  analytical approximations. Additionally, tidal effects usually act over very long time-scales and therefore approximate theories also allow to speed-up the numerical simulations and to explore the parameter space much more rapidly. Secular perturbation theories based on series expansions have been used for hierarchical triple systems. {\\bfx For low values of the eccentricity, the expansion of the perturbation in series of the eccentricity is very efficient \\citep[e.g.][]{Wu_Goldreich_2002}, but this method is not appropriate for orbits that become very eccentric. An expansion in the  ratio of the semi-major axis $a_\\b / a_\\a $ is then preferred, as in this case exact expressions can be computed for the secular system \\citep[e.g.][]{Laskar_Boue_2010}. }   The development to the second order in  $a_\\b / a_\\a $, called the quadrupole approximation, was used by \\citet{Lidov_1961,Lidov_1962} and \\citet{Kozai_1962} for the restricted inner problem (the outer orbit is unperturbed). In this case, the conservation of the normal component of the angular momentum enables the inner orbit to periodically exchange its eccentricity with inclination (the so-called Lidov-Kozai mechanism). There is, however, another limit case to the massive problem, which is the outer restricted problem (the inner orbit is unperturbed). \\citet{Palacian_etal_2006} have studied this case and discussed the existence and stability of equilibria in the non-averaged system. \\citet{Farago_Laskar_2010} derived a simple model of the outer restricted case and described the possible motions of the bodies. They also looked at the quadrupolar problem of three masses and show how the inner and outer restricted cases are related to the general case.  {\\bfx For planar problems, the series expansions in  $a_\\b / a_\\a $ should be conducted to the octopole order \\citep[e.g.][]{Marchal_1990, Ford_etal_2000,Laskar_Boue_2010}, as the quadrupole approximation fails to reproduce the eccentricity oscillations \\citep[e.g.][]{Lee_Peale_2003}. However, the inclinations of the already known hierarchical systems have not been yet determined, and it can be assumed that high values for the eccentricities may also indicate that their mutual inclinations are large as well \\citep[e.g.][]{Laskar_1997,Chatterjee_etal_2008}. }  As for Mercury, Venus and the majority of the natural satellites in the Solar system, close-in bodies undergo significant tidal interactions, resulting that their spins and orbits are slowly modified. The ultimate stage for tidal evolution is the synchronization of the spin and the circularization of the orbit. {\\bfx Indeed, the observed mean eccentricity for planets and binary stars with $ a_\\b < 0.1 $\\,AU is close to zero within the observational limitations \\citep[e.g.][]{Pont_etal_2011}. } Although tidal effects modify the spin in a much shorter time-scale than they modify the orbit, synchronous rotation can only occur when the eccentricity is very close to zero: the rotation rate tends to be locked with the orbital speed at the periapsis, because tidal effects are stronger when the two bodies are closer to each other.   During the formation process, the orbital eccentricity can  increase  due to gravitational scattering, so that the inner bodies become close enough at periapsis for tidal interactions to occur \\citep[e.g.][]{Nagasawa_etal_2008}. The same gravitational scattering is simultaneously responsible for an increase of the mutual inclination of the orbits \\citep[e.g.][]{Chatterjee_etal_2008}, and the fact that inclined systems exchange its inclination with the inner's orbit eccentricity, results that the dissipation in the eccentricity can be transmitted to the inclination of the orbits, and vice-versa. The most striking example is that the spin and the orbit can be completely misaligned \\citep[e.g.][]{Pont_etal_2009,Triaud_etal_2010}. Previous studies on this subject have been undertaken by \\citet{Eggleton_Kiseleva_2001} for binary stars and by \\citet{Wu_Murray_2003} and \\citet{Fabrycky_Tremaine_2007} for a planet in a wide binary. Despite the success obtained by these works in explaining the observations, they all used the same set of equations, derived by \\citet{Eggleton_Kiseleva_2001}, which is not easy to implement and has been obtained in the frame of the inner restricted quadrupolar approximation. As a consequence, their model cannot be applied to a large number of situations, where the outer orbit cannot be held constant, such as regular planetary systems or planets around close binaries.  In this paper we intend to go deeper into the study of hierarchical three-body systems, where the innermost bodies undergo tidal interactions. We do not make any restrictions on the masses of these bodies, and use the quadrupolar approximation for gravitational interactions with general relativity corrections. {\\bfx Our study is then suitable for  binary star systems, planetary systems, and also for planet-satellite systems. } We also consider in our model the full effect on the spins of the two closest bodies, including the rotational flattening of their figures. This allows us to correctly describe the precession of the spin axis and subsequent capture in Cassini states. We adopt a viscous linear model for tides \\citep{Singer_1968,Mignard_1979}, as it provides simple expressions for the tidal torques for any eccentricity value. Since we are interested in the secular behavior, we average the motion equations over the mean anomalies of the orbits and express them using the vectorial methods developed by \\citet{Boue_Laskar_2006}, \\citet{Correia_2009}, and \\citet{Tremaine_etal_2009}.  In Section\\,\\ref{secmodel} we derive the averaged equations of motion that we use to evolve hierarchical systems by tidal effect. In Section\\,\\ref{secevol} we obtain the secular evolution of the spin and orbital quantities in terms of reference angles and elliptical elements, that are useful and more intuitive to understand the outcomes of the numerical simulations. In Section\\,\\ref{appexo} we apply our model to three distinct situations of extra-solar systems: HD\\,80606, HD\\,98800, and HD\\,11964. Finally, last section is devoted to the conclusions. ",
        "conclusions": "Many multi-planet systems have been reported in hierarchical configurations. For the most part, their mutual inclinations are unknown, but the fact that they exhibit significant values in the eccentricities led to think that the inclinations can also be large. In addition, very often the innermost planet in these systems is very close to the star and undergoes tidal evolution. Here we have studied this particular sub-group of multi-planet systems. Using a very general and simplified averaged vectorial formalism, we have shown that inclined hierarchical planetary systems undergoing tidal dissipation can evolve in many different and sometimes unexpected ways.  We re-analyzed the case of HD\\,80606, a system where the planet migrated inwards by a combined tidal-Kozai mechanism, confirming previous results. We looked at the behavior of a planet (or disc) around the binary stars HD\\,98800\\,B and showed that initial prograde orbits may become retrograde and vice-versa, only because of tidal migration within the binary stars. Finally, we studied the regular 2-planet HD\\,11964 system and showed that tidal dissipation combined with gravitational perturbations may lead to a decrease in the mutual inclination, and a fast circularization of the inner orbit.  We have chosen the above three examples, as they are representative of the diversity of behaviors among inclined hierarchical systems. Many other systems are awaiting to be studied. The fact that we use average equations for both tidal and gravitational effects, makes our method suitable to be applied in long-term studies. It allows to run many simulations for different initial conditions in order to explore the entire phase space and evolutionary scenarios. In particular, it can be very useful to put constraints on the inclinations and dissipation ratios of hierarchical planetary systems. {\\bfx Our study can also be extended to systems of binary stars, and to planet-satellite systems. }"
    },
    "1107/1107.3503_arXiv.txt": {
        "abstract": "Recent radio pulsar observations have shown that a number of pulsars display interesting long term periodicities in their spin-down rates.  At least some of these pulsars also undergo sharp changes in pulse profile.  This has been convincingly attributed to the stars abruptly switching between two different magnetospheric states.  The sharpness of these transitions has been taken as evidence against free precession as the mechanism behind the long term variations.  We argue that such  a conclusion is premature.  By performing a simple best-fit analysis to the data, we show that the relationship between the observed spin and modulation periods is of approximately the correct form to be accounted for by the free precession of a population of neutron stars with strained crusts, the level of strain being similar in all of the stars, and consistent with the star retaining a memory of a former faster rotation rate.  We also provide an argument as to why abrupt magnetospheric changes can occur in precessing stars, and how such changes would serve to magnify the effect of precession in the timing data, making the observation of the precession more likely in those stars where such switching occurs.  We describe how future observations could further test the precession hypothesis advanced here. ",
        "introduction": "\\label{sect:intro}  Radio pulsars are remarkably stable rotators.  However, as was apparent from soon after the discovery of the first pulsar, in addition to a smooth secular spin-down, they exhibit other interesting timing features, including sudden spin-ups known as glitches, periodic variations in spin-down rate possibly attributable to free pression, and low-level random timing fluctuations known as timing noise.  The study of all these forms of spin-down variation are of great interest as probes of neutron star structure.  The glitch and precession phenomena are believed to depend sensitively on the coupling between the solid and superfluid phases of the star (see e.g. \\citet{lel99}), while the origin of the timing noise is less well understood, and may be caused by process internal to the star or external (i.e. in the pulsar magnetosphere \\citep{cg81}).   The reported existence of sustained long-period free precession is of particular interest, as  such motions are inconsistent with the pinned superfluid model that has become the standard explanation of large glitches \\citep{ja01, le01,  link03, link06}.  Additionally, timing noise is expected to limit the performance of pulsar timing arrays, whose aim is to detect low frequency (nano-Hertz) gravitational waves \\citep{jetal05}.  As pointed out recently, if timing noise could be modelled and removed from the timing solution, the sensitivity of pulsar timing arrays could be improved \\citep{letal10}.  There are therefore both theoretical and practical motivations for understanding better the origin of timing noise.  As summarised below, recent radio observations have shown that some pulsars with harmonic features in their timing residuals undergo sudden sharp changes in pulse profile which, in one case at least, correlates with sharp changes in spin-down torque \\citep{ketal06, letal10}.  As was convincingly argued, this is good evidence for the star switching back and forth between two different magnetospheric states.  The sharpness of this switching has been taken as evidence that free precession is not a viable mechanism to explain the harmonic timing data \\citep{letal10}, thereby removing the theoretical difficulty discussed above.  However, as we argue below, we believe it is premature to abandon the precession hypothesis as an explanation for the harmonic timing data.  We give two main arguments to support this.  Firstly, the ratios $P/P_{\\rm mod}$ of the spin periods $P$ (of order a second) to the long term periodicities $P_{\\rm mod}$ (of order a few years)  are of the right magnitude to be accounted for by free precession of a star whose crustal stress is due to it retaining a memory of a former faster rotation rate;  few (if any) other known processes act on such long timescales.  We find that the implied breaking strains must be at least as large as $5 \\times 10^{-4}$, with minimum birth spin frequencies of $50$ Hz if only the crust participates in the precession; these numbers increase to $5 \\times 10^{-2}$ and $500$ Hz if the whole star precesses as one.  That free precession of a strained star, with these sorts of minimum strains and birth frequencies, is a viable explanation for the timing behaviour of PSR B1828-11 has already been demonstrated in detail by \\citet{cul03}; what we add here is the demonstration that the scaling of $P_{\\rm mod}$ with respect to $P$ extends over a whole sample of potentially precessing pulsars, consistent with a similar level of strain in each.  Secondly, we argue the sharp magnetospheric switching and precession are not mutually exclusive.  Our argument is qualitative, but, we believe, persuasive.  In brief, suppose some pulsar is delicately balanced between two magnetospheric states.  This would be the case if the energy available to accelerate particles is close to some threshold for some avalanche-like process, say pair production.  Then, if the star is precessing, at some precessional phases the critical energy for the process is exceeded, favouring one magnetospheric state, while at other precessional phases the critical energy is not exceeded, favouring the other magnetospheric state.  That some pulsars should be so delicately balanced is plausible, given the random nulling and mode changing observed even in non-precessing pulsars.  The precession simply adds a bias to the probability of the magnetosphere being in one state or the other.  The corresponding changes in spin-down torque would then serve to amplify the effect of precession on the timing data: it would then be no surprise that switching behaviour and precession have been observed to go hand-in-hand.  It is this precession hypothesis that we explore in this paper, whose structure  is as follows.  In section \\ref{sect:observations} we give a summary of the relevant observations.  In section  \\ref{sect:basic_model} we set out our model.  In section  \\ref{sect:tests} we describe some consistency tests that we have applied to our model.  In section  \\ref{sect:discussion} we discuss our results, and comment on some alternative explanations.  Finally, in section  \\ref{sect:summary} we summarise our findings and give some concluding remarks.  Where relevant, we have assumed neutron star masses of $M = 1.4 M_\\odot$, radii $R = 10^6$ cm, and moments of inertia $I = 10^{45}$ g cm$^2$.  We have made use of the ATNF pulsar database, as described in \\citet{metal05}\\footnote{\\texttt{http://www.atnf.csiro.au/research/pulsar/psrcat}}. ",
        "conclusions": "\\label{sect:summary}  We have presented a model of pulsar timing where the harmonic structure of timing residuals seen in \\citet{letal10} is caused by free precession.  We argued that the relationship between the modulation period and spin period was of the correct form to be explained by free precession of an elastically strained star, the strain being due to the star retaining  a memory of a former faster spin rate \\citep{cul03}.  The details depend upon the fraction of the star that participates in the precession.  In the case where only the crust participates ($I_{\\rm C}/I_\\star \\approx 10^{-2}$) we found the level of strain to be $u \\approx 5 \\times 10^{-4}$, and the stars must have been born spinning as least as rapidly as $\\approx 50$ Hz.  In the case where the whole star participates in the precession  ($I_{\\rm C}/I_\\star \\approx 1$) we found the level of strain to be $u \\approx 5 \\times 10^{-2}$, and the stars must have been born spinning as least as rapidly as $\\approx 500$ Hz.  In either case, the roughly identical levels of strain in the different pulsars might then reflect the universal breaking strain that a crust can withstand.  Alternatively, the nearly identical strain levels may reflect nearly identical spin rates at birth.  We argued that the sharp switching in pulse profile that occurs in at least some of the candidate precessors  is a consequence of the star being delicately balanced between two different magnetospheric states,  with one state being statistically favoured over the other depending upon the precession phase.  That the link between magnetospheric state and precessional phase in only statistical was, we argued, consistent with the completely random mode-switching seen in some non-precessing stars.  The sharpness of the switching reflects the energy of particles at some point in the magnetosphere exceeding some critical threshold for an avalanche-like process, probably pair-production.  Within this model, the magnetosphere variations, locked in phase with the precession, modulate the spin-down torque,  serving to make free precession more readily detectable in such delicately balanced stars as compared to the rest of the population.  It is therefore no surprise that magnetospheric switching has been seen in precession candidates.  The need for the star to be delicately balanced between two different magnetospheric states may play a role in explaining the rarity of the precession phenomenon in the pulsar population, although the precession excitation mechanism (perhaps glitching) may play a role too.  We noted in passing that, in the case where the whole star participates in the precession,  the required asymmetries in the moment of inertia tensor to explain the precession ($\\epsilon_{31} \\sim 5 \\times 10^{-9}$) are similar to the upper limits on asymmetry $\\epsilon_{21}$ that come from assuming the gravitational wave driven spin-down of the millisecond pulsar population.   This may offer hope to gravitational wave astronomers that these upper limits may in fact be estimates, making the targeting of the MSPs by future advanced detectors all the more interesting.  However, if only the crust participated in the precession, the corresponding estimates of $\\epsilon_{31}$ are reduced by two orders of magnitude, comfortably below the upper limits on $\\epsilon_{21}$, removing this seeming coincidence.  The analysis of crustal strain in precession candidates presented here improves upon that of previous works \\citep{ja01, cul03} by considering a larger set of pulsars, as supplied by \\citet{letal10}, and also by imposing the quality cut that at least three precession cycles should fit into the data set for the precession claim to be taken seriously; we suggest that future investigators use a similar quality cut.  Also, the analysis here improves on that of \\citet{ja01} by using the accurate value of the rigidity parameter $b$ calculated by \\citet{cul03}.  There are a number of objections to the model presented here.  Much of our argument is qualitative, rather than quantitative.  The only quantitative piece of evidence we have is the expected value of the modulation period for the slowly spinning stars discussed here, and its scaling with spin frequency, as illustrated in Figure \\ref{fig:P_mod_v_P}.  However, the spread in spin periods (and modulation periods) was not large; more data points are needed to make this convincing.  There is also a serious theoretical problem.  As pointed out by \\citet{ja01} and \\citet{le01}, the superfluid pinning that is believed to be responsible for the larger Vela-like pulsar glitches would drastically reduce the modulation period, eliminating precession as a possible explanation of the observed periodic timing data.  As has been argued by \\citet{link03, link06}, if precession is indeed the mechanism responsible for the timing variations, it would be necessary for any superfluid parts of the star to be completely separated from the charged component, or at least that component that participates in the precession.  This conflicts with microphysical modelling of neutron star interiors, which indicates that the interior superfluid is likely to coexist with a sold crust and, deeper within the star, with magnetic flux tubes.  There is a caveat to be attached to this argument.  The argument against the starquake model of Vela-like glitches is as follows.  The strain relieved at each glitch is much larger than the strain build up by spin-down between glitches, so that the strain cannot possibly be replenished between glitches, i.e. a steady-state solution does not exist.  However,  the large crustal breaking strains  suggested by recent molecular dynamics solutions \\citep{hk09} suggest a possible loop-hole.  A typical Vela glitch releases a strain of order $\\Delta \\nu / \\nu \\sim 10^{-6}$.  The Vela glitches approximately once every $5$ or so years, and is approximately $10^4$ years old.  This corresponds to a total of $\\sim 10^3$ glitches over its lifetime, releasing a total strain of $10^{-3}$, two orders of magnitude smaller that the estimated breaking strain.  So, the Vela and similar pulsars may simply be in the process of divesting themselves of crustal strain built up by spinning down from a much faster rotation rate that applied earlier in their lives.  It follows that providing one is willing    for the Vela and similar glitching pulsars to be in such a non-steady state situation, their large glitches can have a non-superfluid explanation.  However, we can see no good reason why only a small fraction of the total strain should be relieved in each cracking event.  To sum up, we feel that the precession model advanced here is plausible, and gives a sensible account of the observed timing data.  However, the relatively modest number of pulsars with observations of periodic variation in spin-down, and their rather small spread in spin frequency, make forming a firm conclusion impossible.   A larger sample of pulsars with harmonic timing residuals is needed, spanning a wider range of spin and modulation periods.  By placing more points on the modulation period verses spin period plot (Figure \\ref{fig:theory_P-P_mod}),  it should be possible to discriminate between the precession-based model proposed here and the handful of others that have appeared in the literature.  Also, a better picture of the stability of the clock that lies behind the timing data would be invaluable: modulo any irreversible changes in stellar shape, precession should provide high frequency stability; a clear demonstration of a lack of modulation stability would argue against the precessional hypothesis."
    },
    "1107/1107.1395_arXiv.txt": {
        "abstract": "Weak gravitational lensing is a valuable probe of galaxy formation and cosmology. Here we quantify the effects of using photometric redshifts (\\photoz) in galaxy-galaxy lensing, for both sources and lenses, both for the immediate goal of using galaxies with \\photoz\\ as lenses in the Sloan Digital Sky Survey (SDSS) and as a demonstration of methodology for large, upcoming weak lensing surveys that will by necessity be dominated by lens samples with \\photoz.  We calculate the bias in the lensing mass calibration as well as consequences for absolute magnitude (i.e., $k$-corrections) and stellar mass estimates, for a large sample of SDSS Data Release 8 (DR8) galaxies. The redshifts are obtained with the template based \\photoz\\ code ZEBRA on the SDSS DR8 $ugriz$ photometry.  We assemble and characterise the calibration samples ($\\sim$9k spectroscopic redshifts from four surveys) to obtain photometric redshift errors and lensing biases corresponding to our full SDSS DR8 lens and source catalogues. Our tests of the calibration sample also highlight the impact of observing conditions in the imaging survey when the spectroscopic calibration covers a small fraction of its footprint; atypical imaging conditions in calibration fields can lead to incorrect conclusions regarding the \\photoz\\ of the full survey.  For the SDSS DR8 catalogue, we find $\\photozsigma=0.096$ and $0.113$ for the lens and source catalogues, with flux limits of $r=21$ and $r=21.8$, respectively. The \\photoz\\ bias and scatter is a function of \\photoz\\ and template types, which we exploit to apply \\photoz\\ quality cuts.  By using \\photoz\\ rather than spectroscopy for lenses, dim blue galaxies and $L_*$ galaxies up to $z\\sim 0.4$ can be used as lenses, thus expanding into unexplored areas of parameter space. We also explore the systematic uncertainty in the lensing signal calibration when using source \\photoz, and both lens and source \\photoz; given the size of existing training samples, we can constrain the lensing signal calibration (and therefore the normalization of the surface mass density) to within 2 and 4 per cent, respectively. ",
        "introduction": "The current $\\Lambda$CDM cosmological model is dominated by the unknown dark components of the universe: dark matter and dark energy \\citep[e.g.,][]{komatsu_etal:2011}. Gravitational lensing, the deflection of light from distant source galaxies by intervening masses along the line-of-sight \\citep[e.g.,][]{bartelmann_schneider:2001, refregier:2003}, has emerged as an enormously powerful astrophysical and cosmological probe. It is not only sensitive to dark energy through both cosmological distance measures and large-scale structure growth \\citep{albrecht_etal:2006}, but is also sensitive to all forms of matter, including dark matter. Measurements of the statistical distortion of galaxy shapes (or weak gravitational lensing) due to mass along the line-of-sight have been used in numerous studies to constrain the cosmological model, the theory of gravity, and the connection between galaxies and dark matter \\citep[e.g., most recently,][]{mandelbaum_etal:2006, fu_etal:2008, reyes_etal:2010, schrabback_etal:2010}. As a result, many more surveys are planned for the next two decades with weak lensing as a major science driver: KIDS\\footnote{\\texttt{http://www.astro-wise.org/projects/KIDS/}}, DES\\footnote{\\texttt{https://www.darkenergysurvey.org/}}, HSC\\footnote{\\texttt{http://oir.asiaa.sinica.edu.tw/hsc.php}}, Pan-STARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu/public/}}, LSST\\footnote{\\texttt{http://www.lsst.org/lsst}}, Euclid\\footnote{\\texttt{http://sci.esa.int/science-e/www/area/index.cfm?fareaid=102}}, and WFIRST\\footnote{\\texttt{http://wfirst.gsfc.nasa.gov/}}.  In order to fully access the information encoded in gravitational lensing, redshift information is essential, as the conversion from distortions (gravitational shear) to mass depends on the lens and source redshifts via the critical surface density, expressed as \\begin{equation}\\label{E:sigmacrit} \\Sigma_{c}=\\frac{c^2}{4\\pi G} \\frac{D_S}{(1+z_L)^2 D_L D_{LS}} \\end{equation} in comoving coordinates (in physical coordinates, the expression lacks the factor of $(1+z_L)^{-2}$).  Here, $D_L$ and $D_S$ are angular diameter distances to the lens and source, and $D_{LS}$ is the angular diameter distance between the lens and source. Spectroscopic redshifts provide the best accuracy in determining $\\Sigma_c$, but obtaining them for large, statistical samples of both lenses and sources is prohibitively expensive in terms of observing time and instrumentation. As a result, upcoming surveys will rely heavily upon less accurate photometric redshifts (\\photoz) derived from multi-band imaging. Since we require source samples at higher redshift, and as the data pushes into the region of dim galaxies with poor photometry, the \\photoz\\ may worsen even more.   Consequently, it is of immediate interest to quantify how limited redshift accuracy (due to the use of photometric redshifts) propagates into lensing results.  Galaxy-galaxy lensing has been used in the past to quantify the connection between lens galaxies or clusters and their dark matter (DM) halos, in particular the total (average) mass profile around galaxies on $>20$kpc scales and the DM halo occupation statistics \\citep{2005ApJ...635...73H,2006MNRAS.371L..60H,mandelbaum_etal:2006}; can constrain the dark matter power spectrum when used in combination with the galaxy 2-point correlation function \\citep{yoo_etal:2006,cacciato_etal:2009,baldauf_etal:2010,2011PhRvD..83b3008O}; and can constrain the theory of gravity when combined with clustering and redshift-space distortions \\citep{2007PhRvL..99n1302Z,reyes_etal:2010}. This paper addresses the use of photometric redshifts for both sources and lenses in galaxy-galaxy lensing, in the context of the Sloan Digital Sky Survey Data Release 8 \\citep[and references therein; SDSS DR8 hereafter]{eisenstein_etal:prep}. Previous galaxy-galaxy lensing studies with SDSS have been limited to lenses with spectroscopic redshifts \\citep[e.g., ][]{2004AJ....127.2544S,mandelbaum_etal:2006} or photo-$z$ with atypically high accuracy, such as those for Brightest Cluster Galaxies \\citep[e.g., ][]{sheldon_etal:2009}.  Galaxy-galaxy lensing studies with other surveys have typically either involved an unusually large number of passbands yielding exceptionally good \\photoz\\ \\citep[e.g., ][]{kleinheinrich_etal:2006,leauthaud_etal:2011}, or have had limited area coverage and therefore relatively poor statistics compared with the SDSS studies \\citep[e.g., ][]{parker_etal:2007}.  Additional work must be done to allow galaxy-galaxy lensing to achieve its full potential with large, upcoming imaging surveys, and to extend to lens galaxy samples that lack spectra in SDSS. Lens galaxy samples that lack spectra in SDSS tend to be smaller, dimmer galaxies, or galaxies at higher redshifts: it would be interesting to extend galaxy-galaxy lensing studies to their DM halo masses and environments, including mass and redshift dependence.  While studies with other surveys have extended into these regimes, they have typically involved deep but very narrow space-based data with significant cosmic variance \\citep{2006MNRAS.371L..60H,leauthaud_etal:2011} or wider but still relatively noisy ground-based survey data \\citep{2005ApJ...635...73H}.  In order to address galaxy-galaxy lensing based on SDSS DR8 \\photoz, we have calculated a new set of photometric redshifts for the full flux-limited galaxy sample with extinction-corrected model $r<21.8$. We applied the publicly available \\photoz\\ code Z\\\"{u}rich Extragalactic Bayesian Redshift Analyzer (ZEBRA\\footnote{\\tt http://www.exp-astro.phys.ethz.ch/ZEBRA/}, \\citealt{feldmann_etal:2006}) to the SDSS $ugriz$ photometry.  In this work, we demonstrate the \\photoz\\ accuracy by comparing against several spectroscopic samples.  In the course of this process, we identify several concerns related to the observing conditions in the imaging survey in regions that overlap the calibration survey, which will be generally applicable to any upcoming survey.  We define the criterion for a ``better'' \\photoz\\ as a \\photoz\\ that gives minimal scatter in the distribution of $\\zspec-\\zphot$, and hence the lowest scatter in the biases of any redshift-derived quantities, such as the lensing signal or the absolute magnitude. Note that we do not aim for the lowest bias in \\photoz, according to the philosophy that biases can be corrected with a representative calibration sample, but rather the lowest {\\em uncertainty} in the bias\\footnote{Since the scatter defines the uncertainty, we need a representative calibration sample that is large enough to calibrate it. For deeper surveys for which such a large representative sample may not necessarily be available, the definition of 'ideal photometric redshifts' may differ, with a preference for those that show little structure in the \\photoz\\ error as a function of redshift, luminosity, or type.}. This criterion results in a low uncertainty in the actual science analysis. A brief comparison to other publicly available \\photoz\\ catalogues is presented. Improvements in the characterisation of \\photoz\\ for the photometric galaxies can provide a substantial boost in statistics for current studies of the large scale structure. In addition, working with the \\photoz\\ to the survey limiting magnitude with SDSS photometry will provide an ideal test case for future large surveys that rely upon \\photoz\\ for all lensing calculations (which are, necessarily, dominated by galaxies near the flux limit).  With the improved and extended SDSS \\photoz's, we consider the application to galaxy-galaxy lensing. Photo-$z$ for weak lensing source galaxies in SDSS was investigated by \\citealt{mandelbaum_etal:2008}.  Here, we address some additional nuances in the procedures from that work to define a fair calibration sample and to use it to estimate the lensing calibration bias.  We then extend that formalism to use  \\photoz\\ for {\\em lenses}, thus increasing the possible sample of lens galaxies by a factor of $\\sim 40$ over the flux-limited MAIN spectroscopic sample ($r<17.77$) and colour-selected LRGs (luminous red galaxies; $r<19.5$). With an eye to using \\photoz\\ for galaxy-galaxy lensing, we then test how the bias and scatter in the photometric redshifts for the SDSS photometric galaxy sample affect various derived quantities (including the lensing signal calibration, and the estimated luminosities and stellar masses) by direct comparison to their spectroscopic redshift (\\specz) counterparts. While \\cite{kleinheinrich_etal:2005} identify the need for lens redshift information over simply using a lens redshift distribution, our work is the first detailed demonstration of how to calibrate the effect of the lens \\photoz\\ errors on other lensing observables.  The outline of the paper is as follows: we describe the data and calibration subsets in Sec.~\\ref{sec:data}, and test the imaging quality in the regions overlapping the calibration samples for consistency with the average survey quality in Sec.~\\ref{sec:test_calib}. This latter step is crucial for ensuring that measurements using the calibration set are representative of the full SDSS DR8 sample of interest. We justify our choice of \\photoz\\ method in Sec.~\\ref{sec:photoz_method}. The photometric redshift accuracy results (bias and scatter) are discussed in Sec.~\\ref{sec:photoz_results}. From the photometry and the \\photoz, we also derive estimated absolute luminosity and stellar mass, described in Sec.~\\ref{ssec:absmag_bias} and Sec.~\\ref{ssec:stellar_mass}, respectively. Sec.~\\ref{sec:application} then describes the resulting biases of derived quantities for galaxy-galaxy lensing applications. The lensing signal calibration is discussed separately in Sec.~\\ref{sec:calib_lens}, and conclusions are presented in Sec.~\\ref{sec:conclusion}. We use a flat concordance WMAP7 \\citep{komatsu_etal:2011} cosmology ($\\Omega_m=0.27$, $h=0.702$) to calculate luminosity distances $D_L$ and angular diameter distances $D_A$ from redshifts. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have sought to address technical issues required to use photometric redshifts for both lenses and sources in galaxy-galaxy lensing analyses without incurring significant systematic error.  To this end, we have obtained photometric redshifts for SDSS DR8 5-band photometry, using the template-based \\photoz\\ code {\\tt ZEBRA}. Multiple readily available and self-generated \\photoz\\ catalogues were perused and compared to optimize the \\photoz\\ according to the key science goals of the catalogue, galaxy-galaxy lensing, which led us to choose a method with minimal \\photoz\\ scatter for galaxies with $\\zphot>0.1$ (which dominate such analyses).  The underlying assumption is that with a calibration sample of sufficient size to characterize the bias and scatter, any \\photoz\\ biases can simply be calibrated out, so they should not figure into the choice of \\photoz\\ method (provided that they are not so pathological as to make the problem of calibrating them out with a reasonable size calibration sample intractable).  To measure the \\photoz\\ biases and scatter, a calibration set was required. Tests of the calibration set, and corrections for exceptional observing conditions, were required to ensure that its \\photoz\\ scatters and biases accurately represent those of the full SDSS DR8 lensing and source samples. The calibration sample was drawn from several effectively flux-limited spectroscopic samples (six for the shallower lens sample, five for the deeper source sample) consisting of $\\sim$9k galaxies overlapping the SDSS footprint.  The measured \\photoz\\ bias and scatter allowed estimates of their contributions to bias in observed galaxy-galaxy lensing signals.  We considered source and lens \\photoz\\ separately, with emphasis on the newer application to SDSS, lens \\photoz. The use of \\photoz's for source redshifts (in addition to lens spectroscopic redshifts) allows for gravitational lensing signal calibration that is more precise than using an assumed redshift distribution, and allows for selection of specific source subsamples (e.g., brighter or fainter, or with some lens-source separation).  For a known lens redshift, we estimate that the lensing tangential shear can be converted to surface mass density with an accuracy of 2 per cent. Using \\photoz's for lenses will make lenses available at higher redshifts (out to $z\\sim0.5$, limited to red galaxies) and dimmer magnitudes (down to $r\\sim21$ or $\\Mr\\sim-19$, limited to blue), an increase of factors of at least ten compared to SDSS spectroscopic samples.  Rather than lens samples with several tens of thousands of galaxies per stellar mass or luminosity bin as in \\cite{mandelbaum_etal:2006}, the sample sizes will be of order a few $\\times 10^{5}$, which will allow for very high $S/N$ measurements (for which precise calibration of systematics is, therefore, a high priority). The use of lens \\photoz\\ converts to surface mass density with an accuracy of 4 per cent when the lens redshift is $0.1<\\zphot<0.5$.  Besides the lensing signal bias originating in distance relations, biases in the observable or ``stacking quantities'' (required for galaxy-galaxy lensing) arise because such quantities (absolute luminosity, rest-frame colours, stellar mass) also depend on redshift. These biases can be studied because of the availability of best-fit galaxy SED types when obtaining \\photoz\\ from the template-based method. The SED's provide $k$-corrections, which in turn allow for a uniform definition of absolute magnitude across a span of redshifts, as well as stellar mass estimates.  Because we only have a small number of heterogeneous, narrow-field spectroscopic survey samples to serve as the calibration sample, specific cautions were required. For each, the full depth of the photometric survey must be covered to accurately represent the whole sample.  The wide separation of each of the subsamples compensates for the LSS in the narrow field surveys, such that the LSS can be averaged over to represent the (expected) smooth $\\rmd N/\\rmd z$ of the SDSS photometric sample. We found that the differences in survey conditions (such as seeing and sky noise) of the SDSS photometric sample at the locations of the spectroscopic survey area can bias the magnitude completeness and redshift distributions for objects near the limiting magnitude ($r\\sim21.8$).  These then affected the lensing signal bias, where the dim, high redshift source objects have a smaller bias and are highly weighted.  Note that our findings in this paper supersede those in \\cite{mandelbaum_etal:2008}, which (a) had a calibration sample that was a factor of four smaller, and (b) lacked corrections for the atypical observing conditions in SDSS at the position of the COSMOS survey.  While the bias and scatter in \\photoz\\ or \\photoz-derived quantities can be large, we show how calibrating them against the true redshifts allow us to choose a binning width of various lens characteristics appropriate for the uncertainties, and to estimate the correction for the bias.  These biases can now be used to push galaxy-galaxy lensing into a new regime for SDSS."
    },
    "1107/1107.5048.txt": {
        "abstract": "\\noindent We study theoretical implications of direct dark matter searches in the minimal supersymmetric standard model (MSSM). We assume that no accidental cancellations occur in the spin-independent elastic neutralino-quark scattering cross section, but do not impose any relations among the weak-scale MSSM parameters. We show that direct detection cross section below $10^{-44}$ cm$^2$ requires the lightest supersymmetric particle (LSP) neutralino to be close to either a pure gaugino or pure Higgsino limit, with smaller cross sections correlated with smaller admixture of the subdominant components. The current XENON100 bound rules out essentially all models in which the lightest neutralino has the Higgsino fraction between 0.2 and 0.8. Furthermore, smaller direct detection cross sections correlate with stronger fine-tuning in the electroweak symmetry breaking sector. In the gaugino LSP scenario, the current XENON100 bound already implies some fine-tuning: for example, at least 10\\% tuning is required if the LSP mass is above 70 GeV. In both gaugino and Higgsino LSP scenarios, the direct dark matter searches currently being conducted and designed should lead to a discovery if no accidental cancellations or fine-tuning at a level below 1\\% is present. ",
        "introduction": "The Minimal Supersymmetric Standard Model (MSSM) is a well-motivated and extensively studied extension of the Standard Model (SM) at the electroweak scale. Originally introduced as a way to solve the gauge hierarchy problem of the SM, the model turned out to have many other intriguing features. In particular, the MSSM contains an attractive particle dark matter candidate, the lightest neutralino~\\cite{SUSY_DM,DMreviews}. If such a neutralino is the lightest supersymmetric particle (LSP) and R-parity is conserved, it is stable, and its thermal abundance naturally falls in the range indicated by the observed dark matter density.  Probing the microscopic nature of dark matter requires observing interactions of individual dark matter particles with ordinary matter via scattering or annihilation. An extensive experimental effort toward this goal is currently under way. In particular, direct detection techniques, which attempt to detect collisions between ambient dark matter particles and nuclei in the target, have seen significant advances in sensitivity with experiments such as CDMS~\\cite{cdms}, EDELWEISS~\\cite{edel}, and most recently XENON100~\\cite{x100old,x100}. So far, the results of these searches are null: no evidence for dark matter has been observed.\\footnote{DAMA~\\cite{dama} and CoGeNT~\\cite{cogent} collaborations observed effects inconsistent with known backgrounds, which may be due to dark matter scattering. If so, the MSSM dark matter is disfavored, since it falls under the minimal WIMP framework which does not provide a good simultaneous fit to these experiments and XENON100~\\cite{x100}. At the moment, however, the experimental situation is quite confusing, and we will not take DAMA and CoGeNT into account in this study.} The goal of this paper is to investigate the implications of these results, and improvements that are likely to come in the next few years, for the MSSM, with the assumption that the dark matter is made entirely of neutralinos.  In the standard minimal WIMP framework, which is applicable to the MSSM dark matter throughout the model parameter space, the results of direct detection searches are presented as upper bounds on (or perhaps, in the future, measurements of) the spin-independent neutralino-proton elastic scattering cross section at zero momentum exchange,\\footnote{Converting the actually measured event rates into a cross-section bound or measurement requires a number of assumptions, such as local dark matter density (see e.g.\\cite{dmdensity}) and velocity distributions (see e.g.\\cite{veldist}), isospin symmetry of WIMP-nucleon couplings (see e.g.\\cite{isospin}), etc. We will not discuss the potential uncertainties introduced by the conversion in this paper.} which we will call ``direct detection cross section\" for short. As usual with the MSSM, the theoretical prediction for this cross section depends on a large number of model parameters. The traditional approach is to reduce the number of parameters by assuming relations among them, typically at the high energy scale ({\\it e.g.} mSUGRA/cMSSM). Even then, the cross section predictions vary widely depending on the parameters, and the results are usually presented as scatter plots resulting from scanning the parameters within some broad ranges. While such plots provide a useful target for experiments, much potentially interesting information is missing. A typical scatter plot shows the cross section varying over several orders of magnitude, and it is not clear what features of the model correspond to points with higher, or lower, cross sections. Can any qualitative statements about the MSSM be made given the cross section bound of $10^{-44}$ cm$^2$ (roughly corresponding to the current XENON100 result), or some lower future bound? This will be the main focus of this paper. In particular, we will demonstrate a correlation between the direct detection cross section and the amount of fine-tuning in the electroweak sector: roughly speaking, model points with lower direct detection cross sections are more fine-tuned.  \\begin{figure}[t] \\begin{center} \\centerline { \\includegraphics[width=4.5in]{diagrams.pdf} } \\caption{Feynman diagrams contributing to spin-independent elastic scattering of the neutralino dark matter particle on a nucleon in the MSSM.} \\label{fig:Fdiagrams} \\end{center} \\end{figure}  In order to be as general as possible, we will treat all weak-scale MSSM parameters as independent, without assuming any relations among them. The tree-level processes contributing to the direct detection cross section are shown in Fig.~\\ref{fig:Fdiagrams}. The key assumption underlying our analysis is that {\\it no accidental cancellations take place} among various contributions to the direct detection cross section in the MSSM. By ``accidental\", we mean a cancellation which is exact only on a measure-zero hypersurface inside the full MSSM parameter space. Equivalently, an accidental cancellation is indicated by an anomalous sensitivity of the cross section to the MSSM parameters (measured, for example, by its logarithmic derivative) along at least one direction in the parameter space. In particular, any cancellation between the $s$- and $t$-channel diagrams in Fig.~\\ref{fig:Fdiagrams} would be accidental, since they depend on different sets of MSSM parameters.\\footnote{Of course, different MSSM parameters may be related once the SUSY-breaking sector is understood, so that a cancellation that appears accidental from the weak-scale point of view may in fact be natural in the full theory. Such a situation, however, appears extremely unlikely in the particular situations where we apply the ``no accidental cancellation rule\" in this study. For example, a cancellation between the $s$- and $t$-channel diagrams in Fig.~\\ref{fig:Fdiagrams}  would require a complicated relation involving squark and gaugino soft masses, the $\\mu$ parameter, $\\tan\\beta$, and the Higgs mass terms. It is very difficult to imagine a SUSY-breaking model producing such a relation.} Thus, for making qualitative statements, it is sufficient to consider only one of the diagram classes; the other one will, at worst, produce an order-one correction to the cross section. We will focus on the $t$-channel diagrams, Fig.~\\ref{fig:Fdiagrams} (a). We make this choice because three of the five MSSM parameters which enter these diagrams, $\\mu$, $\\tan\\beta$, and $m_A$, also enter the tree-level prediction for the $Z$ mass. In this way, the direct detection cross section is connected to electroweak symmetry breaking.  When comparing with experimental data, we will assume that the local dark matter density is at its canonical value, $0.3$ GeV/cm$^3$, and that dark matter is made entirely of the MSSM neutralinos. However, the main results of the paper concern theoretical predictions for direct detection cross section and are independent of this assumption. Moreover, in most of the analysis, we will {\\it not} impose the constraint that the neutralino relic density predicted by the standard thermal decoupling calculation be consistent with observations. This is motivated partly by the desire to keep the analysis as general as possible: Statements made without the relic density constraint would be applicable even if there is non-thermal production of neutralinos, late inflation, other deviations from standard FRW cosmological evolution, {\\it etc}. More technically, it is motivated by the fact that the object we focus on, the $t$-channel direct detection cross section, depends on just five MSSM parameters, while the relic density is a function of many more. Imposing this constraint would thus greatly complicate the analysis. (As an exception, we will impose a mild version of the relic density constraint in the part of the analysis dealing with Higgsino dark matter, where this will be required to make an interesting statement about fine-tuning.)  Before we proceed, let us remark on a few related analyses in the literature. The connection between direct detection cross section and fine-tuning in the electroweak sector was noted, in the mSUGRA context, in Ref.~\\cite{GPMM} (see also~\\cite{sa,DG}). A study of this connection in a general MSSM framework similar to ours, focusing on the bino-LSP scenario, appeared in Ref.~\\cite{KN}. A recent study of ``generic\" direct detection cross sections in a general MSSM framework was presented in Ref.~\\cite{FS}. While our approach is similar, our setup is even more general: in particular, we do not require gaugino mass unification, and (for the most part) do not impose the thermal relic density constraint. Electroweak fine-tuning, which did not enter the analysis of~\\cite{FS}, plays the central role in our discussion. We also discuss the impact of the recent XENON100 bound in the general MSSM framework, with the main assumption being the absence of accidental cancellations. This discussion complements Refs.~\\cite{akula,strumia,xenon_msugra} which interpreted the XENON100 result within mSUGRA and other constrained MSSM frameworks.  The rest of the paper is organized as follows. In section~\\ref{sec:setup}, we set the notation and collect the main formulas used in the analysis. The main results of the paper are contained in section~\\ref{sec:partial}, where we present a set of scatter plots demonstrating correlations between the direct detection cross section and physical quantities such as Higgsino fraction of the neutralino and the amount of fine-tuning in the electroweak symmetry breaking. These plots assume real, positive MSSM parameters. This assumption is lifted in section~\\ref{sec:full}, where negative and complex soft masses are considered, and it is shown that the interesting correlations persist once the points with accidental cancellations in the cross section are eliminated. Finally, we conclude in section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper, we considered the theoretical implications of direct dark matter searches in the MSSM framework. Unlike almost all previous analyses, we did not require any relations among the MSSM parameters, such as mSUGRA, and did not impose thermal relic density constraints. Instead, we make the MSSM parameter space tractable by assuming that no accidental cancellations take place among the $s$-channel and $t$-channel contributions to the spin-independent elastic neutralino-quark cross section, and focusing on the $t$-channel diagrams. These only depend on 5 MSSM parameters. We performed extensive scans over these parameters, in order to identify correlations between the direct detection cross sections and other quantities of physical importance. The following simple picture emerges from our analysis:  \\begin{itemize}  \\item If the LSP is a generic mixture of Higgsino and gauginos with order-one admixture of each component, the direct detection cross section is always above $2\\times 10^{-44}$ cm$^2$. Such models are already severely constrained: for example, essentially all models with a Higgsino fraction between 0.2 and 0.8 are ruled out by XENON100.  \\item Lowering direct detection cross section below the $10^{-44}$ cm$^2$ level requires that the LSP be either pure gaugino (bino or wino), or pure Higgsino. In both cases, smaller cross sections correlate with higher purity ({\\it i.e.}, smaller admixture of the subdominant components) of the LSP.  \\item If the LSP is a gaugino, smaller direct detection cross sections correlate with stronger fine-tuning in the electroweak symmetry breaking sector, since they require higher values of the $\\mu$ parameter. The current XENON100 bound already implies non-trivial fine-tuning, albeit rather mild at this point: for example, 1 part in 10 or worse tuning is required if the LSP mass is above 70 GeV. Future experiments could put stronger stress on the model: for example, if no signal is seen at the proposed XENON100 upgrade, fine-tuning of 1\\% or worse would be required for LSP mass above 100 GeV. This would present a new ``little hierarchy problem\" for the MSSM (at least if the idea that the neutralino makes up all of the dark matter is taken seriously).  \\item If the LSP is a Higgsino, no correlation between direct detection cross section and fine-tuning can be established. However, if an additional mild assumption, the one-sided thermal relic density constraint (see sec.~\\ref{sec:higgsino}) is imposed, all points with direct detection cross section below about  $2\\times 10^{-44}$ cm$^2$ have an LSP mass of 1 TeV or above, and electroweak fine-tuning at the level of $0.25$\\% or worse. Non-observation of a signal at the proposed XENON100 upgrade would imply this level of tuning in the Higgsino LSP scenario.  \\end{itemize}  Once again, it is worth emphasizing that these conclusions are very general, and apply in the full ``phenomenological\" MSSM, without undue theoretical prejudice~\\cite{no_prej}. The degree of fine-tuning in the EWSB is a widely accepted quantitative ``figure of merit\" used to assess relative theoretical attractiveness of various regions of the MSSM parameter space. We established a clear correlation between this measure and the direct detection cross section. The main conclusion of our analysis is that, if the MSSM is the true model of microscopic physics and dark matter, and no fine-tuning at a sub-percent level is present, the direct dark matter searches currently being conducted and designed should lead to a discovery.  A similar analysis could be performed in other models of electroweak symmetry breaking with particle dark matter candidates. An interesting direction for future work is the so-called next-to-minimal supersymmetric standard model (NMSSM), an MSSM with an additional singlet field in the Higgs sector. This model significantly alleviates the fine-tuning due to the Higgs mass lower bound, and the additional ``singlino\" admixture may be present in the LSP, affecting the direct detection cross section. It would be interesting to see if the correlations discussed here persist in the NMSSM.  \\vskip0.8cm \\noindent{\\large \\bf Acknowledgments} \\vskip0.3cm This research is supported by the U.S. National Science Foundation through grant PHY-0757868 and CAREER grant No. PHY-0844667.  This research was  supported in part by the National Science Foundation under Grant No. PHY05-51164. MP is grateful to Michael Peskin and to Aaron Pierce for useful remarks. MP acknowledges the hospitality of the Kavli Institute for Theoretical Physics at Santa Barbara, where part of this work has been completed. BS acknowledges the hospitality of the Theoretical Advanced Study Institute (TASI-11) at the Univeristy of Colorado at Boulder.  \\begin{appendix}"
    },
    "1107/1107.1440_arXiv.txt": {
        "abstract": "In general neutron stars in binaries are spinning. Due to the existence of millisecond pulsars we know that these spins can be substantial.  We argue that spins with periods on the order a few dozen milliseconds could influence the late inspiral and merger dynamics.  Thus numerical simulations of the last few orbits and the merger should start from initial conditions that allow for arbitrary spins.  We discuss quasi-equilibrium approximations one can make in the construction of binary neutron star initial data with spins.  Using these approximations we are able to derive two new matter equations.  As in the case of irrotational neutron star binaries one of these equations is algebraic and the other elliptic.  If these new matter equations are solved together with the equations for the metric variables following the Wilson-Mathews or conformal thin sandwich approach one can construct neutron star initial data.  The spin of each star is described by a rotational velocity that can be chosen freely so that one can create stars in arbitrary rotation states.  Our new matter equations reduce to the well known limits of both corotating and irrotational neutron star binaries. ",
        "introduction": "Several gravitational wave detectors such as LIGO~\\cite{LIGO:2007kva,LIGO_web}, Virgo~\\cite{VIRGO_FAcernese_etal2008,VIRGO_web} or GEO~\\cite{GEO_web} have been operating over the last few years, while several others are in the planning or construction phase~\\cite{Schutz99}. One of the most promising sources for these detectors are the inspirals and mergers of binary neutron stars. In order to make predictions about the last few orbits and the merger of such systems, fully non-linear numerical simulations of the Einstein Equations are required. To start such simulations we need initial data that describe the binary a few orbits before merger. The emission of gravitational waves tends to circularize the orbits~\\cite{Peters:1963ux,Peters:1964}. Thus, during the inspiral, we expect the two neutron stars to be in quasi-circular orbits around each other with a radius that shrinks on a timescale much larger than the orbital timescale. This means that the initial data should have an approximate helical Killing vector $\\xi^{\\mu}$. To incorporate these ideas we will use the Wilson-Mathews approach~\\cite{Wilson95,Wilson:1996ty}, which is also known as conformal thin sandwich formalism~\\cite{York99}, for the metric variables. The Wilson-Mathews approach has already been successfully used by several groups together with matter equations describing the neutron stars in either corotating~\\cite{Baumgarte:1997xi,Baumgarte:1997eg, Mathews:1997pf,Marronetti:1998xv,Tichy:2009yr} or irrotational~\\cite{Bonazzola:1998yq,Gourgoulhon:2000nn,Marronetti:1999ya, Uryu:1999uu,Marronetti:2003gk,Taniguchi:2002ns,Taniguchi:2003hx, Uryu:2005vv,Uryu:2009ye} states. There have also been attempts to include intermediate rotation states~\\cite{Marronetti:2003gk,Baumgarte:2009fw}. However, as we will discuss in more detail later, these approaches have certain drawbacks, because they do not correctly solve the Euler equation for the fluid. Thus, so far there is no canonical formalism to describe neutron star binaries with arbitrary spins. As pointed out by Bildsten and Cutler~\\cite{Bildsten92}, the two neutron stars cannot be tidal locked, because the viscosity of neutron star matter is too low. Hence barring other effects like magnetic dipole radiation the spin of each star remains approximately constant. This means that initial data sequences of corotating configurations for different separations cannot be used to approximate the inspiral of two neutron stars. On the other hand, sequences of irrotational configurations can be used to approximate the inspiral of two neutron stars without spin. This fact explains why irrotational initial data are far more popular today. Nevertheless, astrophysical neutron stars will have a non-zero spin. Therefore a corotating configuration at some particular separation does have its place as a possible initial configuration with spin. It will just not remain corotating during the subsequent time evolution. Of course real neutron stars will likely have spins that have periods different from the orbital period, and the spin direction may not be aligned with the orbital angular momentum. Thus it would be highly desirable to have a formalism that can be used to generate initial data for arbitrary initial spins.  In order to judge how important spins might be let us discuss a few order of magnitude estimates. A typical neutron star has a mass of about 1.4 solar masses ($M_{\\odot}$) and a radius on the order of 15km. From Kepler's law the orbital period \\begin{equation} P_o \\sim  \\left(\\frac{d}{50\\mbox{km}}\\right)^{3/2} \\left(\\frac{M_{\\odot}}{M}\\right)^{1/2}  6\\mbox{ms} \\end{equation} is on the order of a few milliseconds during the last orbit before merger where the separation $d\\sim 50\\mbox{km}$. Thus systems with spin periods that are much larger than $P_o$ should be treatable as approximately irrotational, while systems with spin periods of a few milliseconds (such as millisecond pulsars) cannot be regarded as irrotational. Another way of judging how important spins could be during the evolution is to look at the dimensionless spin magnitude. If we assume that the spin $S$ of a neutron star with mass $m$ and radius $R$ is related to its spin period $P$ by $S= I (2\\pi/P)$ with $I\\sim m R^2$ we find that the dimensionless spin has a magnitude of \\begin{equation} \\frac{S}{m^2} \\sim \\left(\\frac{R}{15\\mbox{km}}\\right)^2 \\frac{M_{\\odot}}{m} \\frac{3 \\mbox{ms}}{P} . \\end{equation} Thus millisecond pulsars have a dimensionless spin of order one. As in the case of binary black holes~\\cite{Campanelli:2006uy, Campanelli:2006fy,Campanelli:2007ew, Tichy:2007hk,Marronetti:2007wz,Tichy:2008du}, spins of this magnitudes could have a significant influence on the merger dynamics. This means that neutron stars with spin periods of a few dozen milliseconds or less should not be considered irrotational. One could of course imagine that the neutron stars spin down before they enter the strongly relativistic regime of the last few orbits before merger that is usually considered in numerical relativity simulations. In order to address this question let us look at the famous double pulsar PSR J0737-3039 which is the only neutron star binary where both spin periods and spin down rates are known~\\cite{Lyne:2004cj}. Star A has mass $m_A=1.34 M_{\\odot}$ and spin period $P_A=23\\mbox{ms}$, while star B has $m_B=1.25 M_{\\odot}$ and $P_B=2.8\\mbox{s}$. The orbital period is $P_o=2.4\\mbox{h}$. From these numbers one derives that the system will merge in about $85\\mbox{My}$ due to the emission of gravitational waves. Both stars are currently spinning down at a rate of $\\dot{P}_A=1.7\\times 10^{-18}$ and $\\dot{P}_B=8.8\\times 10^{-16}$~\\cite{Lyne:2004cj}. If one assumes that this spin down is due to magnetic dipole radiation and defines the characteristic ages given by $\\tau_A=P_{A}/(2\\dot{P}_{A})= 210\\mbox{My}$ and $\\tau_B=P_{B}/(2\\dot{P}_{B})= 50\\mbox{My}$, one finds that the spin period of each star obeys~\\cite{Shapiro83} \\begin{equation} P_{A/B}(t) = P_{A/B}(0) \\sqrt{1+\\frac{t}{\\tau_{A/B}}} , \\end{equation} where the time $t=0$ is the time today. From this it is clear that the periods at merger (at $t=85\\mbox{My}$) will be $P_{A}(t)=27\\mbox{ms}$ and $P_{B}(t)=4.6\\mbox{s}$. Thus star A will not spin down enough to be well approximated by an irrotational configuration by this time. This example shows that neutron stars in binary systems can have appreciable spins a few orbits before merger. Of course since only about ten binary neutron stars have been observed so far~\\cite{Lorimer:2005bw} it is not clear yet how common binary neutron stars with high spins are. However, since there are numerous millisecond pulsars it seems reasonable to expect that neutron stars in binaries can also have millisecond spin periods. Hence the widely held belief that only irrotational configurations are realistic is not necessarily correct. It is thus necessary to develop initial data for binary neutron stars with arbitrary spins. In the next sections we will describe what approximation one can make to derive a formalism that allows for this possibility. We will see that our new equations reduce to well known accepted results in both the corotating and irrotational cases.      Throughout we will use units where $G=c=1$. Latin indices such as $i$ run from 1 to 3 and denote spatial indices, while Greek indices such as $\\mu$ run from 0 to 3 and denote spacetime indices. The paper is organized as follows. Sec.~\\ref{BNSequations} lists the General Relativistic equations that govern binary neutron stars described by perfect fluids. We use three approximate quasi-equilibrium conditions to simplify these equations. We find two new matter equations that allow as to set up binary neutron stars with arbitrary spins. In Sec.~\\ref{NewtonLimit} we consider the Newtonian limit of our new equations. We conclude with a discussion of our method in Sec.~\\ref{discussion}. In the appendix we discuss our quasi-equilibrium conditions for a simple case. ",
        "conclusions": "\\label{discussion}  Realistic neutron stars in binaries will be spinning. From observations of millisecond pulsars we know that these spins can be substantial enough to influence the late inspiral and merger dynamics of the binary.  There have been prior attempts to construct initial data for spinning neutron stars. In~\\cite{Marronetti:2003gk} (hereafter MS) the Euler equation is not solved directly. Rather it is replaced by an equation equivalent to $ \\frac{h}{u^0} + ^{(3)}\\!\\tilde{u}_k V^k = -C $. However, as already pointed out by MS, this equation agrees with the integrated Euler Eq.~(\\ref{EulerInt}) only for the corotating and the irrotational case. Thus in general the Euler equation is violated in the MS approach. Furthermore, MS split $u^i/u^0$ and not $^{(3)}\\!\\tilde{u}^i$ into an irrotational and a rotational part (see Eq.~(\\ref{utilde-split})). This has two consequences. First, their equations do not have the correct limit in the irrotational case. And second, since $u^{\\mu}/u^0$ is not a purely spatial vector it is inconsistent to set $u^i/u^0$ equal to something like $D^i \\phi$ which is a purely spatial vector. This explains why the continuity equation of MS has no shift terms unlike in Eq.~(\\ref{continuity4}) and in~\\cite{Shibata98,Teukolsky98}. When MS compare their results for a particular corotating case with~\\cite{Uryu:1999uu} they find that their approach introduces errors of about 2\\% in the angular momentum.  Another approach to include spin that is aligned with the orbital angular momentum was proposed in~\\cite{Baumgarte:2009fw}, hereafter BS.  This approach does not seek to analytically integrate the Euler equation as we have done here.  Instead the divergence of Eq.~(\\ref{Euler2}) is set to zero, which leads to another elliptic equation.  However, as pointed out first by Gourgoulhon~\\cite{Baumgarte:2009fwErr}, in general the Euler equation itself is not satisfied if we only enforce its divergence to be zero.  Hence the BS approach can lead to initial data that do not obey the Euler equation.  Furthermore, the boundary condition given by BS for their new elliptic equation seems to imply that the star surface is always at the same location. If we consider the usual numerical treatment where we start from an initial guess for the stars which is iteratively refined, it is unclear how the star surface can change during the iterations.  The purpose of this paper is thus to introduce a new method for the computation of binary neutron star initial data with arbitrary rotation states. Our method is derived from the standard matter equations of perfect fluids together with certain quasi-equilibrium assumptions. We assume that there is an approximate helical Killing vector $\\xi^{\\mu}$ and that Lie derivatives of the metric variables with respect to $\\xi^{\\mu}$ vanish. We also assume that scalar matter variables such as $h$ or $\\rho_0$ have Lie derivatives that vanish with respect to $\\xi^{\\mu}$. However, as discussed in appendix \\ref{appendix_assumptions} the Lie derivative of the fluid velocity $u^{\\mu}$ is expected to be non-zero for arbitrary spins. We split the fluid velocity $u^{\\mu}$ into an irrotational piece (derived from a potential $\\phi$) and a rotational piece $w^i$, and assume that only the irrotational piece has a vanishing Lie derivative (see Eq.~(\\ref{assumption1})) with respect to $\\xi^{\\mu}$. This can be interpreted the natural generalization of the irrotational case where one commonly assumes $\\pounds_{\\xi} h u^{\\mu} = 0 $. Furthermore we know that the spin of each star remains approximately constant since the viscosity of the stars is insufficient for tidal coupling~\\cite{Bildsten92}. To incorporate this fact, we use Eq.~(\\ref{assumption2}) which is based on the assumption that $w^i$ is constant along the star's motion described by the irrotational velocity piece $\\nabla^{\\mu} \\phi$. Since $\\nabla^{\\mu} \\phi$ is equivalent to the velocity of the star center, this latter assumption captures the fact that the spin or rotational velocity $w^i$ of each star remains approximately constant. With these two assumptions the Euler equation simplifies to Eq.~(\\ref{Euler3}). In order to analytically integrate Eq.~(\\ref{Euler3}) we use the additional assumption (\\ref{assumption3}) that $\\tilde{u}^0$ and $\\gamma_{ij}$ are constant along the field lines of the rotational velocity piece. We then arrive at the two matter equations Eq.~(\\ref{continuity4}) and (\\ref{h_from_Euler}). These equations reduce to well known equations~\\cite{Shibata98,Teukolsky98} for the irrotational case of $w_i=0$. They also reduce to the corotating limit (where $V^i=0$) as is evident from Eqs.~(\\ref{continuity3}) and (\\ref{EulerInt}) which are written in terms of $V^i$. Furthermore, our equations reduce to the correct Newtonian limit.  The elliptic equation in Eq.~(\\ref{continuity4}) can to solved (for $\\phi$) together with the Eqs.~(\\ref{ham_mom_dtK0}) for the metric variables once the enthalpy $h$ is known.  However, the enthalpy given by Eq.~(\\ref{h_from_Euler}) depends on the metric variables, $\\phi$ and $w^i$. Apart from their dependence on $w^i$ this set of equations has a similar structure as for the case of irrotational neutron stars (where $w^i$ vanishes). The standard way (see e.g. ~\\cite{Tichy:2009yr}) to solve such a mixture of elliptic and algebraic equations is by iteration, where at each step we first solve the elliptic equations for a given $h$ and then use the algebraic Eq.~(\\ref{h_from_Euler}) to update $h$. At each step we also need to specify $w^i$. One way to do this would be by choosing a constant $\\bar{w}^i$ as in Eq.~(\\ref{wbar_omega_cross_r}). Note however, that other choices for $w^i$ are possible. We plan to investigate these possibilities in future numerical studies of our new method. For such studies it might useful to use a numerical code like LORENE~\\cite{Grandclement-etal-2000:multi-domain-spectral-method, Gourgoulhon:2000nn,Grandclement02,lorene_web} or SGRID~\\cite{Tichy:2006qn,Tichy:2009yr,Tichy:2009zr}, where the star surface is always at a domain boundary so that the boundary condition in Eq.~(\\ref{starBC}) can be easily implemented."
    },
    "1107/1107.3997_arXiv.txt": {
        "abstract": "We use more than 4,500 microflares from the Reuven Ramaty High Energy Solar Spectroscopic Imager ({\\it RHESSI}) microflare data set (Christe et al., 2008, Ap.~J., 677, 1385) to estimate electron densities and volumetric filling factors of microflare loops using a cooling time analysis.  We show that if the filling factor is assumed to be unity, the calculated conductive cooling times are much shorter than the observed flare decay times, which in turn are much shorter than the calculated radiative cooling times.  This is likely unphysical, but the contradiction can be resolved by assuming the radiative and conductive cooling times are comparable, which is valid when the flare loop temperature is a maximum and when external heating can be ignored.  We find that resultant radiative and conductive cooling times are comparable to observed decay times, which has been used as an assumption in some previous studies.  The inferred electron densities have a mean value of $10^{11.6} \\ {\\rm cm}^{-3}$ and filling factors have a mean of $10^{-3.7}$.  The filling factors are lower and densities are higher than previous estimates for large flares, but are similar to those found for two microflares by Moore et al.~(Ap.~J., 526, 505, 1999). ",
        "introduction": "Energy release in flares in solar active regions occurs over many orders of magnitude, from large flares to microflares with as low as a millionth the energy content of large flares.  The latter are A and B {\\it GOES} class events, which occur more frequently than large flares with a negative power law distribution in number of flares as a function of energy release extending over many decades in energy \\citep{Lin84,Dennis85,Crosby93,Feldman97,Wheatland00,Nita02, Paczuski05}, which suggests a common energy release mechanism. There are some features in common among flares of all sizes, such as radiation in multiple wavelength bands and similar X-ray light curves [see, {\\it e.g.}, \\citet{Fletcher10} for a review].  However, observational differences also exist between small and large flares. Large flares are associated with higher temperatures than small flares \\citep{Feldman95,Caspi10}.  Also, weaker hard X-ray flares may have steeper spectra than more energetic ones \\citep{Christe08,Hannah08}.  Statistical studies of hard X-ray microflares have become more comprehensive since the launch of the Reuven Ramaty High Energy Solar Spectroscopic Imager ({\\it RHESSI}) satellite \\citep{Lin02}.  {\\it RHESSI} achieves a lower energy cutoff in the X-ray spectrum than previous detectors.  By using removable shutters, {\\it RHESSI} allows observation of both large and small flares [see, {\\it e.g.}, \\citet{Hannah10} for a description].  A recent study \\citep{Christe08} used a new flare-finding technique to identify over 24,000 microflares from March 2002 - March 2007.  Statistical analyses of the microflare properties were carried out by \\citet{Christe08} and \\citet{Hannah08}. This large data set allows for unprecedented studies of flare properties.  Here, we use this large data set to infer the electron density and the volumetric filling factor of the microflare loops in the {\\it RHESSI} data set.  The volumetric filling factor is the fraction of the flare loop volume from which radiation is detected.  While the filling factor is often thought of as a robust parameter for flare loops, it should be noted that its determination is potentially instrument- and resolution-dependent.  A previous estimate of the density assumed the filling factor was unity \\citep{Hannah08}.  We use a cooling-time analysis [see, {\\it e.g.}, \\citet{Moore80}], to argue that the volumetric filling factors of microflare loops may be considerably smaller than unity, implying densities considerably higher than the estimates from \\citet{Hannah08}.  The filling factors and densities we find are consistent with a previous study of two microflares observed with {\\it Yohkoh} \\citep{Moore99}.  The {\\it RHESSI} microflare data set is described in Sec.~\\ref{sec-observations}.  The analytical technique is reviewed and critiqued in Sec.~\\ref{sec-technique}.  Results and uncertainty estimates are presented in Sec.~\\ref{sec-results}.  Finally, conclusions are discussed in Sec.~\\ref{sec-disc}. ",
        "conclusions": "\\label{sec-disc}  In this paper, {\\it RHESSI} microflare data are used to estimate the volumetric filling factor and the electron density of microflare loops using an analysis of cooling times.  If the filling factor is assumed to be unity, then the conductive cooling time of the loop is much smaller than the observed decay time, which itself is much smaller than the radiative decay time.  This is difficult to justify physically.  Alternately, if one invokes the hypothesis that the radiative and conductive cooling times are comparable at the moment when the flare temperature passes through its maximum value (and that cooling due to expansion and flare heating are negligible at that time), one can solve for the filling factor and density.  Mean values for the whole distribution are $\\phi \\sim 10^{-3.7}$ and $n_{e} \\sim 10^{11.6} \\ {\\rm cm}^{-3}$.  Our weakest assumption is that flare heating stops at the peak time of hard X-rays.  Since the hard X-ray time profile is a convolution of heating and cooling, heating does not necessarily stop at the hard X-ray peak time.  If heating is present during the decay, even at a low level, the cooling times could be longer than derived for the case without heating, and the filling factor could be larger than derived here.  Our estimate of mean densities are higher than those reported by \\citet{Hannah08}.  The authors are aware of only one other systematic study of microflare loop densities or filling factors \\citep{Moore99}, who used {\\it Yohkoh} to study two microflare strands.  Using an identical analytical technique as the present study, they found densities between $n_{e} \\sim 10^{10} \\ {\\rm cm}^{-3}$ and $10^{11.6} \\ {\\rm cm}^{-3}$, with filling factors between $10^{-3.2}$ and $10^{-2.8}$.  Thus, the values in the two microflares studied by \\citep{Moore99} from {\\it Yohkoh} are in good agreement with results from the large {\\it RHESSI} data set studied here.  We now compare the present results with observations of large flares. \\citet{Culhane94} found $\\phi \\sim 1$ and $n_{e} \\sim 3 \\times 10^{11} \\ {\\rm cm}^{-3}$ for an M-class flare, \\citet{Varady00} found $\\phi \\sim 0.01-0.2$ and $n_{e} \\sim 7 \\times 10^{9} - 1.5 \\times 10^{10} \\ {\\rm cm}^{-3}$ for a C-class flare, \\citet{Aschwanden01b} found $\\phi \\sim 1$ and $n_{e} \\sim 1.5 \\times 10^{10} \\ {\\rm cm}^{-3}$ for the Bastille Day flare, \\citet{Teriaca06} found $\\phi \\sim 0.2-0.5$ and $n_{e} \\sim 10^{10} \\ {\\rm cm}^{-3}$ for a C-class flare, and \\citet{Raymond07} found $\\phi \\geq 0.01$ and $n_{e} \\sim 10^{11} \\ {\\rm cm}^{-3}$ for X-class flares.  Other examples of filling factors include 0.3 in active region loops \\citep{Landi09}, 0.04-0.07 in coronal holes \\citep{Abramenko09}, and near unity in many coronal hole jets \\citep{Doschek10} but 0.03 in another study \\citep{Chifor08}.  As noted earlier, the determination of filling factors can depend on detector resolution and wavelength.  Nonetheless, the filling factors obtained here for microflares are at least 10 times smaller than those reported for large flares, which is likely statistically significant. Densities are slightly higher for microflares in the present study than for larger flares in previous studies.  Given the characteristics of the microflare data set considered here compared to large flares, it is perhaps not surprising that the filling factors are small.  The microflare loops in the present study occur in active regions, just as large flares do.  The mean sizes of the loops in the present study are comparable to those of larger flares.  The loops have energies of $n_{e} k_{B} T V \\sim 10^{28} {\\rm \\ ergs}$ deposited into them by the flare (using characteristic values from Fig.~\\ref{fig-rawdata}), which is about $10^{4}$ times smaller than large flares.  Thus, the loops are of similar size but acquire less energy, which could lead to a smaller filling factor.  To estimate the size of the region for which radiation is detected, we note the radiating volume is $V_{*} = \\phi V$.  Assuming that the length of the radiating plasma is $L$, the effective width $w_{*}$ of the radiating plasma is given by $V_{*} \\sim \\pi (w_{*} / 2)^{2} L$. For the present parameters, this implies loops with total thickness $w_{*} \\sim \\phi^{1/2} w \\simeq w / 100 \\simeq 4 \\times 10^{6} \\ {\\rm cm}$, which would imply there is an unseen substructure of thin strands within the flare loops.  Various lines of evidence indicate that there are smaller-scale structures in the corona, e.g.~\\citet{Mullan90}.  Radio polarization data point to the existence of structures in the corona which are $\\sim ~100$ km in size \\citep{Melrose75}.  The possibility that 100 km structures are associated with collapsing magnetic reconnection sites in the corona was discussed by \\citet{Mullan80}: using constraints on the collapse time-scales and coronal Alfv\\'en speeds, transverse dimensions of order 100 km were found to be typical of reconnection sites in the corona.  More recently, there is evidence that X-class flare loops are composed of thin threads from high resolution observations, with structure at scales of a few arcseconds ($1'' \\simeq 10^{8} {\\rm \\ cm}$) and below \\citep{Dennis09,Kontar10, Krucker10}.  There is also abundant evidence from the footpoints of flaring loops that most of the emission is spatially unresolved, such as in TRACE white-light flares \\citep{Hudson06}.  \\citet{Xu06} found a core region within a halo region in two X-class white-light flares, reporting a ratio of the area of the core to the halo of 4\\% and 25\\%, respectively.  Also, simulations of loops comprised of many small scale filaments were able to reproduce cooling characteristics of large flare loops observed with TRACE \\citep{Warren03}.  Thus, the conclusion that there is small substructure of flare loops is not without precedent.  The prediction of small scale loops has implications for the heating mechanism of the flare loops.  \\citet{Hannah08} estimated that the non-thermal power in accelerated electrons during the time of peak emission in the {\\it RHESSI} microflares is $10^{26} \\ {\\rm erg \\ s}^{-1}$.  For loops of area $10^{14} {\\rm \\ cm}^{2}$ as predicted by the present results, the energy deposition rate per unit area would be $10^{12} {\\rm \\ erg \\ s}^{-1} {\\rm \\ cm}^{-2}$.  This is an enormous value, as discussed in \\citet{Krucker10}.  Hence, if the filling factor is indeed $\\sim 10^{-4}$, then microflares are unlikely heated by electron beams.  A recent model that the flare energy is transported by Alfv\\'en waves \\citep{Fletcher08} would not be ruled out by the data.  An interesting result of the present study is that the conductive and radiative cooling times derived by assuming their equality at the time of maximum temperature are comparable to the observed microflare decay times.  A possible ramification of this result is that it lends credence to the assumption that the conductive and radiative times are comparable to the decay time, $\\tau_{R} \\sim \\tau_{C} \\sim \\tau_{D}$. This is relevant to stellar flare studies in which plasma parameters were obtained under such assumptions \\citep{Haisch83,Stern83}.  In addition, a previous study of stellar flares \\citep{Mullan06} found results of this model are largely consistent with independently determined plasma parameters.  If one believes that the scaling analysis cooling times actually represent physical cooling times for loops, the result suggests that conductive and radiative cooling act at comparable levels to cool flare loops, at least for the microflares in the present study.  Future work could include efforts to incorporate physical effects left out of the model as summarized at the end of Sec.~\\ref{sec-technique}. Also, future studies could further try to minimize the systematic errors discussed in Sec.~\\ref{sec-results}.  These can be addressed both with observations and with numerical modeling.  Also, the study of the cooling times of individual events will help determine the validity of the cooling time analysis.  The authors thank G.~Holman and B.~Dennis for helpful conversations. RNB and PAC gratefully acknowledge support by NSF grant PHY-0902479, NASA's EPSCoR Research Infrastructure Development Program and the West Virginia University Faculty Senate Research Grant program.  DJM is supported in part by the Delaware Space Grant.  IGH is supported by a STFC rolling grant and by the European Commission through the SOLAIRE Network (MTRN-CT-2006-035484).  RPL is supported in part by the WCU grant (No. R31-10016) funded by the Korean Ministry of Education, Science and Technology.  HSH, RPL, and SK are supported through NASA contract NAS 5-98033 for {\\it RHESSI}."
    },
    "1107/1107.3445_arXiv.txt": {
        "abstract": "We work out the impact that the recently determined time-dependent component of the Pioneer Anomaly (PA), if interpreted as an additional exotic acceleration of gravitational origin with respect to the well known PA-like constant one, may have on the orbital motions of some planets of the solar system. By assuming that it points towards the Sun, it turns out that  both the semi-major axis $a$ and the eccentricity $e$ of the orbit of a test particle would experience secular variations. For Saturn and Uranus, for which modern data records cover at least one full orbital revolution, such predicted anomalies  are up to $2-3$ orders of magnitude larger than the present-day accuracies in empirical determinations of their orbital parameters from the usual orbit determination procedures in which the PA was not modeled. Given the predicted huge sizes of such hypothetical signatures, it is unlikely that their absence from the presently available processed data can be attributable to an \\virg{absorption} for them  in the estimated parameters caused by the fact that they were not explicitly modeled. The magnitude of a constant PA-type acceleration  at $9.5$ au cannot be larger than  $9\\times 10^{-15}$ m s$^{-2}$ according to the latest observational results for the perihelion precession of Saturn. ",
        "introduction": "According to the latest analysis\\cite{Tury011} of extended data records of the Pioneer $10/11$ spacecraft, the small frequency drift\\cite{And98,And02} (blue-shift) observed analyzing the navigational data of both the spacecraft, known as Pioneer Anomaly (PA),  may present a further time-dependent component in addition to the well known  constant one. Both linear and  exponential models were proposed\\cite{Tury011} for the PA; according to the authors of Ref.~\\refcite{Tury011}, the exponential one is directly connected to  non-gravitational effects\\cite{Toth09} since it takes into account the possible role of the on-board power generators suffering a radioactive decay.  In this letter we  work out the orbital effects of such a new term in the hypothesis that the time-dependent PA component is due to some sort of  long-range modification of the known laws of gravitation resulting in an additional anomalous acceleration with respect to the nearly sunward constant one, having magnitude\\cite{And02} \\eqi \\left|A_{\\rm Pio}\\right|=(8.74\\pm 1.33)\\times 10^{-10}\\ {\\rm m\\ s^{-2}},\\eqf in terms of which the constant part of the PA has often been interpreted. Indeed, in this case it should act on the major bodies of the solar system as well, especially those whose orbits lie in the  regions in which the PA manifested itself in its presently known form. In this respect, we will not consider the exponential model. Recent studies\\cite{Berto08,Riv09,Berto010,Riv010a,Riv010b,Fran011,Riv011}, partly preceding the one in Ref.~\\refcite{Tury011}, pointed towards a mundane explanation of a large part of the PA in terms of non-gravitational effects pertaining the spacecraft themselves. ",
        "conclusions": "If  the time-dependent linear part of the PA were a genuine effect of gravitational origin causing an extra-acceleration,  it would affect the orbital motions of the planets with secular variations of their semi-major axis and eccentricity if it were  oriented towards the Sun. The magnitude of such putative effects for Saturn and Uranus is up to $2-3$ orders of magnitude larger than the current accuracy in determining their orbital motions from the observations without modeling the PA itself. A partial removal of a PA-like signature may, in principle, have occurred in the standard parameter estimation procedure, so that the entire planetary data set should be re-processed with ad-hoc modified dynamical models explicitly including a hypothetical time-dependent PA-type acceleration as well.  However, supported by the results of analogous tests for the constant form of the PA actually implemented by independent teams of astronomers with different purposely modified dynamical models, we feel that it is unlikely that such a potential removal of a PA-type signature  can really justify the absence of such a huge effect in the currently available processed planetary data. Concerning the existence of a constant PA-type acceleration at heliocentric distances of about $9.5$ au as large as $1.6\\times 10^{-9}$ m s$^{-2}$, latest results about the maximum allowed size of anomalies in the  perihelion precession of Saturn yield a figure 5 orders of magnitude smaller."
    },
    "1107/1107.2055_arXiv.txt": {
        "abstract": "Recent results of the Pierre Auger (Auger) fluorescence detectors indicate an increasingly heavy composition of ultra-high energy (UHE) cosmic rays (CRs). Assuming that this trend continues up to the highest energies observed by the Auger surface detectors we derive the constraints this places on the local source distribution of UHE CR nuclei. Utilizing an analytic description of UHE CR propagation we derive the expected spectra and composition for a wide range of source emission spectra. We find that sources of intermediate-to-heavy nuclei are consistent with the observed spectra and composition data above the ankle. This consistency requires the presence of nearby sources within $60$~Mpc and $80$~Mpc for silicon and iron only sources, respectively. The necessity of these local sources becomes even more compelling in the presence nano-Gauss local extragalactic magnetic fields. ",
        "introduction": "Ultra-high energy CRs with energies above $10^{19}$~eV arrive at Earth with a frequency of less than one event per square-kilometer-year ({\\it i.e.}~with an energy flux of 30~eV~cm$^{-2}$~s$^{-1}$) in $\\pi$ steradian. The recently completed Pierre Auger observatory \\cite{Abraham:2004dt} with a surface detector covering an area of about $3000$~km$^{2}$ is thus able to detect up to several hundred UHE CR events per year. However, due to the steeply falling CR spectrum the arrival frequency drops by about two orders of magnitude as we go up in energy by one decade, leaving only a few events per year detectable at energies around $10^{20}$~eV. Hence, the statistical uncertainty associated with the upper end of the spectrum is still large, limiting our knowledge of UHE CR composition and origin.  Before their arrival, UHE CRs must propagate across the astronomical distance between their source and Earth. The relevant interactions of UHE CR nuclei during propagation are Bethe-Heitler pair production and photo-disintegration in collisions with photons of the cosmic background radiation. The cross-sections of both these processes rise quickly above the threshold values of about 1~MeV and 10~MeV, respectively, in the nuclei's rest-frame. At even higher energies above 150~MeV pion production turns on and becomes the dominant energy loss process. However, for the UHE CR cutoff energies and composition we consider, which are motivated by the most recent Auger results, this process never plays a dominant role and may be safely neglected.  Provided the propagation time from their sources to Earth is greater than their energy loss time, UHE CRs invariably undergo these energy loss interactions. The break-up of nuclei via photo-disintegration produces lower mass nuclei with the same Lorentz factor. For heavy nuclei, the most dominant transitions are one-nucleon and two-nucleon losses. Secondary heavy nuclei remain close to the photo-disintegration resonance and quickly disintegrate further to lighter nuclei. Hence, on resonance, the initial mass composition of the sources is quickly shifted to lower atomic number values and in general shows a strong dependence on CR energy.  The arriving flux from an ensemble of UHE CR sources can be expected to contain suppression features at the high energies at which the photo-disintegration processes turn on and the nuclei particle's attenuation length decreases. Most prominently, for the case of a proton-dominated spectrum, the flux is expected to be suppressed by the {\\it Greisen-Zatsepin-Kuz'min} (GZK) cutoff~\\cite{Greisen:1966jv,Zatsepin:1966jv} due to resonant pion photo-production interactions with the cosmic microwave background (CMB). Intriguingly, a suppression of the CR spectrum at about $5\\times10^{19}$~eV has been observed at a statistically significant level~\\cite{Abbasi:2007sv,Abraham:2008ru}. However, since a similar feature may also appear in the spectrum from nuclei primaries, the observation of such a feature provides little clue as to the underlying composition. Such a suppression feature becomes a cutoff in the arriving flux if the attenuation length drops below the distance to the nearest UHE CR source. Thus, the shape of the suppression/cutoff feature does contain information about the source distribution, as has been investigated already for the case of UHE CR protons~\\cite{Aharonian:1994nn,Berezinsky:2002nc,Taylor:2008jz}.  On their arrival at Earth, UHE CRs interact with molecules in the atmosphere and deposit their energy in the form of extensive air showers. The characteristics of these showers along the shower depth $X$ (in g/cm${^2}$) contain vital UHE CR composition information. On average, proton-induced showers reach their maximum development, $\\langle X_{\\rm max} \\rangle$, deeper in the atmosphere than do showers of the same energy generated by heavier nuclei. Accompanying this effect, the shower to shower fluctuation of $X_{\\rm max}$ about the mean, $\\mathrm{RMS}(X_{\\rm max})$, is larger for proton-induced showers than for iron-induced showers of the same energy. As a result, measurements of both $\\langle X_{\\rm max} \\rangle$ and $\\mathrm{RMS}(X_{\\rm max})$ can be used to infer the average chemical composition of the UHE CRs as a function of energy.  Auger has now released their first measurement results of $\\mathrm{RMS}(X_{\\rm max})$~\\cite{Unger:2011ry,Abraham:2010yv}, along with those of $\\langle X_{\\rm max} \\rangle$. These results seem to imply that the UHE CR spectrum contains a large fraction of intermediate or heavy mass nuclei, becoming increasingly heavy at high energies ($\\sim 10^{19.5}$~eV). Furthermore, the small values of $\\mathrm{RMS}(X_{\\rm max})$ measured by Auger also imply that UHE CRs are composed of species with a relatively narrow distribution of charge at the highest measured energies, containing little or no protons or light nuclei. In this way, the new $\\mathrm{RMS}(X_{\\rm max})$ measurements not only confirm and reinforce the conclusions drawn from their earlier average depth of shower maximum measurements, but also provide complementary information that enables one to constrain the distribution of the various chemical species present within the UHE CR spectrum. Curiously, recent corresponding measurements by the Telescope Array UHE CR detector do not appear to agree with the Auger results \\cite{Matthews:2011zz}, and instead seem to be consistent with a light composition as was indicated by its predecessor HiRes~\\cite{Abbasi:2009nf}. Though an understanding of this conflict is crucial for future progress in the field, we here chose to adopt only the Auger results.  Interestingly, the intermediate-heavy nuclei scenario emerging from these measurements may find consistency with the presently observed mild departure from isotropy and absence of point sources observed at UHE \\cite{:2010zzj}. This could occur through the introduction of $\\sim 10^{\\circ}$ deflections in the turbulent ($\\mu$G) Galactic magnetic field structure \\cite{Giacinti:2011uj} or in 0.1~nG turbulent extragalactic field structure \\cite{Dolag:2004kp}, from an anisotropic distribution of nearby sources. Such large deflections would limit the prospects for future UHE CR astronomy for all but the brightest and closest sources. The correlation of UHE CR with nearby objects, however, still awaits to be resolved within an intermediate-heavy nuclei framework, and deserves further consideration.  Following our previous investigations into the composition of UHE CRs~\\cite{Hooper:2009fd}, an intermediate-heavy (silicon-to-iron) type composition was found to be motivated by the present complete Auger data set, with a hard ($\\alpha< 2$) injection spectral index \\footnote{This spectral index refers to an injection spectrum of the form ${\\rm d}N/{\\rm d}E\\propto E^{-\\alpha}$.}. \\setcounter{pub}{\\value{footnote}} Furthermore, this agreement received only very mild improvements by the additional consideration of an admixture of species. Building further on these results, we here focus on the local UHE CR source distribution inferred to exist if the observed trend in the composition continues up to highest energies observed by the ground array ($\\sim 10^{20.2}$~eV). By considering simple silicon and iron source composition scenarios, we investigate how far away these sources can afford to be without detrimentally damaging the agreement with the Auger data. To aid this investigation we utilise an analytic description of UHE CR nuclei propagation already developed by the authors~\\cite{Ahlers:2010ty}. ",
        "conclusions": "\\label{Summary}  Following the motivation that UHE CRs consist of heavy nuclei, whose sources have been suggested to be local~\\cite{Wibig:2007pf,Piran:2010yg}, we have here looked closer at the requirements on their source distribution. In this work we obtained concrete quantitative constraints on the UHE CR source population. Making explicit use of the recently provided Auger spectral and shower composition results, along with detailed UHE CR nuclei modeling, we investigated whether consistency may be found with this data using single source composition models.  For the case of negligible extragalactic magnetic fields, we have demonstrated that a simplified analytic description agrees well with the spectrum and composition results obtained from the complete Monte Carlo description. Utilising this analytic description, we made a scan over the source spectral index and exponential energy cutoff, for a single source composition scenario, to obtain the goodness-of-fit contours. Taking into account the systematic errors in the flux and composition measurements, we found that hard ($\\alpha<2$) source spectral indices and intermediate cutoff energies ($E_{\\rm Fe, max}\\sim 10^{20.5}-10^{21}$~eV) for intermediate-to-heavy nuclei could provide a good fit to the full set of Auger UHE CR measurements above 10$^{19}$~eV.  Through the consideration of shells of UHE CR sources, we investigated the proximity of the UHE CR nuclei sources required if the presently observed trend of an increasingly heavy composition continues up to the highest energies observed by the ground array ($10^{20.2}$~eV). By varying the size of the local void up to the nearest source, we investigated how detrimental the effect of this was on the goodness-of-fits contours. Provided that the sources maximum energy lay below $10^{22}$~eV, we found that the nearest sources had to be within $60$~Mpc and $80$~Mpc for silicon and iron only sources respectively.  If, however, extragalactic magnetic fields are sufficiently strong ($>$pG), the arriving flux from the different nuclei source shells is altered considerably. Though this effect is weakest at the highest energies, a local void of sources scenario with such an intervening field, may alter the arriving flux at energies below the cutoff feature, and thus alter the goodness-of-fit contour landscapes obtained in Fig.~\\ref{Contours}. However such ($<$nG strength) fields are found to be unable to alter significantly our upper bound on the nearest source distance.  The requirement for local candidate objects able to satisfy the Hillas criterion \\cite{Hillas:1985is} for $>10^{20}$~eV nuclei, whilst at the same time not disintegrating these nuclei during the acceleration process, places a non-trivial dual condition on the source environment. At first consideration both the required proximity of the sources and their required nuclei tolerant radiation fields would seem to exclude GRBs as viable candidates. However, recent motivations for a Galactic GRB-type origin \\cite{Calvez:2010uh} and the possibility that heavy nuclear species are produced during the GRB explosion \\cite{Metzger:2011xs} demonstrate that the case for a GRB origin remains viable. Due to the dependence of photo-disintegration rates on the radiation field level, acceleration sites far away from the central engine are naturally favoured. In this regards, acceleration within AGN jets or their radio lobes \\cite{O'Sullivan:2009sc} are able to satisfy the above mentioned dual condition.  With regards energetics, the UHE CR source luminosity density required to power the UHE CR population above $10^{19}$~eV is $\\sim 5\\times 10^{44}$~erg~Mpc$^{-3}$~yr$^{-1}$ \\cite{Waxman:1995dg}. For the hard sources spectra motivated in this paper, approximately the same source luminosity density would exist up to the cutoff energy. Thus an average local luminosity per source of $\\sim 10^{43}/N_{\\rm s}$~erg~s$^{-1}$ is required within the local 60~Mpc region, where $N_{\\rm s}$ is the number of contributing sources. Such a luminosity is roughly $1/N_{\\rm s}$ that of the 20-40~keV X-ray luminosity of local AGN \\cite{Beckmann:2005gs,Baumgartner:2010}, of which $\\sim$10 sit within 60~Mpc from Earth. Whether it is reasonable for these local AGN to channel such a large fraction of their non-thermal luminosity into UHE CR flux, and whether a sufficiently enhanced nuclear composition is able to be achieved within their acceleration sites, however, remains unclear \\cite{Pe'er:2009rc,GopalKrishna:2010wp}.  Our results demonstrate that exciting consequences follow from the intermediate-to-heavy nuclei component uncovered by Auger measurements. With nuclei photo-disintegration inevitably occurring during propagation, tough constraints are placed on the source proximity and environment. Furthermore, future Auger spectral and composition measurements are anticipated to soon tighten these constraints. With few candidate sources within the present upper bound distance suggested ($<60$-$80$~Mpc), the puzzle as to the UHE CR origin both remains and becomes even more intriguing."
    },
    "1107/1107.5169_arXiv.txt": {
        "abstract": "A large fraction of the information collected by cosmological surveys is simply discarded to avoid lengthscales which are difficult to model theoretically. We introduce a new technique which enables the extraction of useful information from the bispectrum of galaxies well beyond the conventional limits of perturbation theory.  Our results strongly suggest that this method increases the range of scales where the relation between the bispectrum and power spectrum in tree-level perturbation theory may be applied, from $\\kmax \\sim 0.1 \\hmpc$  to $\\sim 0.7 \\hmpc$. This leads to correspondingly large improvements in the determination of galaxy bias.  Since the clipped matter power spectrum closely follows the linear power spectrum, there is the potential to use this technique to probe the growth rate of linear perturbations and confront theories of modified gravity with observation. ",
        "introduction": "At the largest observable scales, density fluctuations in the Universe have sufficiently low amplitude to allow a simple approximation to model their evolution, in the form of linear perturbation theory. But at progressively smaller scales this formalism breaks down, and more complex prescriptions are required to predict the clustering patterns. While in recent years there has been steady progress improving our understanding of the nonlinear matter power spectrum (e.g. \\cite{1996MNRAS.280L..19P, 2003MNRAS.341.1311S, RPTCrocceScoccimarro}), such developments are unlikely to keep pace with the  precision demanded by the ambitious galaxy surveys of the future. If we are to take full advantage of their capabilities, these surveys  will require knowledge of the matter power spectrum to sub-percent precision at lengthscales well into the nonlinear regime.  The galaxy bispectrum is able to offer insight into key topics in cosmology.  Along with a theoretical model of its evolution, a measurement of the bispectrum enables properties of the galaxy bias to be determined \\cite{2002MNRAS.335..432V}. It also potentially holds a signature of inflation, in the form of primordial non-Gaussianity. However, like the power spectrum, the bispectrum is known to be strongly influenced by nonlinearities. To obtain robust parameter constraints, it is usual to discard information above a limiting $\\kmax$, an obviously undesirable state of affairs. At low redshifts we are restricted to working on extremely large scales,  typically $\\kmax \\lesssim 0.1 \\hmpc$ \\cite{2010Sefusatti}, where cosmic variance severely hampers our precision. This value of $\\kmax$ may be improved with more detailed modelling, going beyond tree-level perturbation theory to one-loop or Renormalised Perturbation Theory \\cite{RPTCrocceScoccimarro, 2010Sefusatti}; see also \\cite{2008PhRvD..78b3523S}. These methods are able to extend the $k$-range of applicability by a moderate amount over tree-level, at the cost of considerable complexity.  An alternative to modelling the full matter power spectrum is to manipulate the real space density field such that the effects of nonlinear growth are suppressed, an approach that has previously been demonstrated to successfully enhance the baryon acoustic oscillations (BAO)  \\cite{EisSeoRecon}. There is, therefore, motivation to study such transformations as a general tool for cosmological analysis.  The simplest class of reconstruction methods are local, monotonic mappings from the initial to final density fields. The Gaussianisation process proposed by Weinberg \\cite{1992MNRAS.254..315W} enforces Gaussianity in the field's one-point distribution, and this has been shown to recover the shape of the linear power spectrum \\cite{Neyrinck2010}. Some other filters, including $\\log (1+\\delta)$ \\cite{Neyrinck2009}, perform comparably well. What such transformations have in common is the penalisation of the highest density regions of the field, a feature we shall exploit further. Our goal is to extend the range of the relatively simple tree-level perturbation theory by manipulating the field in real space. This work begins with the ansatz that the nonlinear behaviour of the density field  is localised not only at high wavenumbers in Fourier space, but also within high density regions in real space. ",
        "conclusions": "Determining the form of the dark matter power spectrum remains a principal goal of modern cosmology. We have demonstrated how the simple prescription of clipping high density peaks allows a much larger volume of $k$-space to be incorporated into such predictions. This leads to a substantially more accurate determination of the bias, and consequently tighter constraints on cosmological parameters.  For the highly evolved field of the Millennium simulation ($\\sigma_8 = 0.9$ at redshift zero), we find that the maximum wavenumber $\\kmax$ may typically be extended by this technique from $\\sim 0.1 \\hmpc$ to beyond $\\sim 0.5 \\hmpc$. The full impact  of this extension upon signal to noise becomes apparent upon noting that the number of triangles contributing to the bispectrum computations scales in proportion to $\\kmax^6$.  In this letter we have chosen to focus on the galaxy bias parameters, but it should be clear that this technique is extensible to quite generic measurements in cosmology. At the heightened precision with which the matter bispectrum can now be studied, it is natural to consider whether improved constraints on $f_{NL}$ are attainable.  Although any signature of primordial non-Gaussianity is expected to be most prominent at the largest scales, the abundance of modes at high $k$ may compensate for the smaller signal. The clipping procedure may also prove valuable in improving cosmological constraints on the neutrino mass \\cite{2008PhRvL.100s1301S, 2009PhRvD..79b3520I, 2009PhRvD..80h3528S, 2010MNRAS.405..168L}, and in modified gravity models where the linear growth rate differs from General Relativity.  There are a number of ways in which the technique could be developed further.  A closer-to-optimal approach is to evaluate the relationship between the preferred truncation threshold and particular wavenumber values. Indeed the ultimate limit on $\\kmax$ is at present uncertain, since larger truncation fractions extend the range of validity beyond $0.7 \\hmpc$. The compensation of the clipping window, and moving to redshift space, where there will undoubtedly be challenges due to `Fingers of God', will be topics of forthcoming work. Exploitation of the bispectrum to the level of precision that will soon be observationally accessible will likely require the incorporation of scale-dependence into the bias model. Considering these future directions, it seems reasonable to conclude that density filtering schemes may establish themselves as an important tool for cosmological analysis.   We thank Peder Norberg, Sylvain de la Torre, Benjamin Joachimi and Thomas Kitching for helpful discussions. FS and AFH acknowledge the Dark Cosmology Centre's visitor programme. FS and CH acknowledge support from the European Research Council under the EC FP7 grant number 240185. JBJ is supported by the Danish National Research Foundation's Sophie \\& Tycho Brahe Programme in Astrophysics."
    },
    "1107/1107.5443_arXiv.txt": {
        "abstract": "{} {We performed observations with the Effelsberg 100-m radio telescope to measure flux densities and polarised emission of sources selected from the ``Deep X-ray Radio Blazar Survey\" (DXRBS) to better define their spectral index behaviour in the radio band, with the aim to construct a homogeneous sample of blazars. } {Sources were observed at four different frequencies with the Effelsberg 100-m telescope. We complemented these measurements with flux density data at 1.4\\,GHz derived from the NRAO VLA Sky Survey (NVSS). } { The spectral indices of a sample of faint blazars were computed making use of almost simultaneous measurements. Sixty-six percent of the sources can be classified as ``bona fide'' blazars. Seven objects show a clearly inverted spectral index. Seventeen sources previously classified as flat spectrum radio quasars (FSRQ) are actually steep spectrum radio quasars (SSRQ). The flux densities obtained with the Effelsberg 100-m telescope at 5\\,GHz are compared with the flux densities listed in the Green Bank (GB6) survey and in the Parkes-MIT-NRAO (PMN) catalogue. About 43\\% of the sources in our sample exhibit flux density variations on temporal scales of 19 or 22 years. } {We confirm that 75 out of 103 sources of the DXRBS are indeed FSRQs. Twenty-seven sources show a spectral index steeper than $-0.5$ and should be classified as SSRQs. Polarised emission was detected for 36 sources at 4.85\\,GHz. The median value of the percentage of polarised emission is ($5.8\\pm0.9$)\\%. Five sources show rotation measure (RM) values $>$ 200 rad m$^{-2}$.  } ",
        "introduction": "Blazars are an extreme class of active galactic nuclei (AGN) characterized by high luminosity, rapid variability, and high polarisation. In the radio band, blazars are core-dominated objects with apparent superluminal speeds along relativistic jets pointed close to the observer's line of sight. Blazars have flat spectral indices and include FSRQs and BL Lacertae objects, the counterparts of high- and low-luminosity radio galaxies. Owing to their orientation with respect to the line of sight, blazars represent less than 5\\% of all AGN, a quite rare class of sources.  The first blazar samples were selected at relatively high limiting flux densities in the radio or X-ray band, $\\sim$1 Jy or a few times 10$^{-13}$ ergs cm$^{-2}$ s$^{-1}$, respectively \\citep{stickel91,wood84,kollgaard96,gioia90,stocke91,perlman96}.  In the last decade of the past century, many efforts have been made to select samples of blazars that could be representative of the whole population \\citep{marcha94,laurent99,caccianiga99,bondi01}. A deeper, more comprehensive sample of blazars has been constructed by \\citet{perlman} and by \\citet{landt}:  the ``Deep X-ray Radio Blazar Survey\" (DXRBS). The DXRBS has been constructed by cross-correlating all ROSAT sources of the WGCAT catalogue \\citep{white95} and radio sources with flat radio spectra. The radio flux densities have been taken from available catalogues, like the Green Bank GB6 \\citep{gregory}, NORTH20CM \\citep{white92}, and Parkes-MIT-NRAO PMN \\citep{griffith}.  The DXRBS sample is currently the faintest, down to $\\sim$50\\,mJy at 5\\,GHz and power $\\sim$10$^{24}$ W\\,Hz$^{-1}$, and most extensive blazar sample with nearly complete optical spectroscopic identifications.  The radio spectral information of the DXRBS sources is based on non-simultaneous measurements at just two frequencies. Simultaneous flux density measurements at several frequencies are needed to obtain more reliable spectral indices.  In this paper we present the results from the simultaneous multi-frequency campaign made with the Effelsberg 100-m telescope. In Section \\ref{sec:observation} we summarise the observations and data processing. In Section \\ref{sec:results} we present results from the observations. Conclusions are presented in sections \\ref{sec:comm-sources} and \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} Our investigation aims at veryfing the spectral index classification of the DXRBS sample. We define a spectral index $\\alpha\\le-0.5$ as ``steep''. If $\\alpha>-0.5$ it is considered ``flat''. However, the spectra often show a complex shape, i.e. not a steep or flat straight line. This makes the classification more difficult. We have attempted a ``crude'' classification source by source in Table~\\ref{tab:spix}, while spectral indices are available in the plots we produced for each source (see Appendix A).  The majority of sources, 66 out of 102 (for one source we do not have enough measurements) can be considered $bona fide$ blazars, namely those with ``flat'', ``steep-flat'', ``inverted'', and possibly also ``flat-steep'' spectra in the present classification.  Several sources changed their spectral index classification with respect to that reported in the original list by \\citet{perlman} and by \\citet{landt}. Among them we have nine sources with a ``convex'' spectral index peaking in the GHz regime, which should be more properly defined as giga-Hertz peaked sources (GPS). These sources cannot be considered blazars. We found that 17 sources previously classified as FSRQs are actually SSRQs. Moreover, we found six sources with spectral indices even steeper than $-0.7$, failing one of the original selection criteria. In total, we found that 27 sources can be classified as SSRQs showing a spectral index $<-0.5$. This result suggests that a statistically meaningful comparison between FSRQs and SSRQs in the same flux density limited complete sample is possible.  We can anticipate that all 42 sources belonging to our sample observed so far with the EVN at 5\\,GHz have been detected (Mantovani et al., in prep.). Six sources, previously classified SSRQs (and found with $\\alpha<-0.7$), host AGN. Two of them show a core-jet structure, while the remaining four are point-like at a resolution better than 5 milli-arcseconds. The EVN to Effelsberg flux density ratio ranges from 0.05 to 0.8. Most of their radio emission is therefore not coming from the cores of these sources.  Fourty-three percent of the DXRBS sample show significant variability on a time scale of 19--22 years; 11\\% show short-term variability comparing Green Bank flux density measurements taken one year apart in 1986 and 1987.  About 25\\% of the sources show polarised emission above the detection limits of our observations. We also extracted polarimetric information from the NVSS to collect at least the three measurements needed to compute the RM diminishing the $n\\pi$ ambiguity. The majority of the sources show $\\mid$RM$\\mid$ $<200$ rad m$^{-2}$ in the source rest frame, possibly produced by the Faraday depth of the Galactic foreground. In nine cases, namely J0029.0$+$0509, J0322.6$-$1335, J0434.3$-$1443, J0510.0$+$1800, J0518.2$+$0624, J0744.8$+$2920, J1400.7$+$0425, J1419.1$+$0603, and J1648.4$+$4104 the rest frame RM are in the range of 200$<\\mid$RM$\\mid<$1940 rad m$^{-2}$. The highest value found is 1938.8 rad m$^{-2}$ for the source J0434.3$-$1443."
    },
    "1107/1107.0099_arXiv.txt": {
        "abstract": "We present long term near-IR light curves for two nearby AGN: NGC~3783 and MR~2251-178. The near-IR data are complemented with optical photometry obtained over the same period of time. The light curves in all bands are highly variable and good correlations can be seen between optical and NIR variations. Cross-correlation analysis for NGC~3783 suggests that some disc near-IR emission is present in the J-band flux, while the H and K-bands are dominated by emission from a torus located at the dust sublimation radius. For MR~2251-178 the cross-correlation analysis and the optical--near-IR flux-flux plots suggest that the near-IR flux is dominated by disc emission.  We model the optical to near-IR Spectral Energy Distributions (SED) of both sources and find that disc flaring might be a necessary modification to the geometry of a thin disc in order to explain the observations. The SED of MR~2251-178 gives some indications for the presence of NIR emission from a torus. Finally, we consider the implications of the standard alpha disc model to explain the different origin of the variable NIR emission in these AGN. ",
        "introduction": "Variability was very early established as a trade mark of Active Galactic Nuclei (AGN). Variability has also been used as a key tool to derive physical properties of AGN: characteristic time scales were used to infer sizes of the emitting regions, lags between the ionising continuum and the line response have been used to determine Black Hole masses ($M$), multi-wavelength light curves have been used to study the geometrical and physical connections between the different regions around the central engine.  The scenario accepted until recently for the interplay of the emitting regions is that variability is driven by the emission from the X-ray corona located close to the central Black Hole (e.g., Collin-Souffrin, 1991; Krolik et al.\\ 1991; Clavel et al., 1992; Collier et al.\\ 1999; Cackett et al.\\ 2007). The negligible optical inter-band lags was early evidence that pointed towards reprocessing of high-energy photons by the accretion disc, where the characteristic distances between the different emitting regions correspond to the light travelling time (Collin-Souffrin, 1991; Krolik et al.\\ 1991; Clavel et al., 1992). Later on, the measurement of short lags between the X-ray emission and the optical, and the leading of the X-rays whenever significant lags were determined (e.g., Edelson et al., 1996; Wanders et al., 1997; Shemmer et al., 2001; Desroches et al., 2006), lent support to this picture (see also Nandra et al., 2000 for evidence supporting reprocessing from the correlation between the UV flux and the spectral shape of the X-ray emission), since for intrinsic disc variability shorter wavelengths should lag longer wavelength emission by long (viscous) time-scales. Also, light curves showed that the amplitude of the X-ray variations was much larger than that seen in the optical, which can be explained by the damping of the signal during the disc reprocessing. However, a full picture of the interplay between the X-ray corona and the disc might not be complete just yet, as new evidence seems to suggest that long term optical variability is driven by accretion rate fluctuations in the disc itself (see Section 3.1 and also Ar\\'evalo et al., 2008, 2009, and references therein).  Long term ($\\sim$ years), well sampled optical, UV, and X-ray monitoring of AGN is available for several sources (e.g., Breedt et al., 2010; Chatterjee et al., 2009; Shemmer et al., 2001; Giveon, et al., 1999; T\\\"urler et al., 1999; Korista et al., 1995; Clavel et al., 1991, Clavel, Wamsteker \\& Glass, 1989). However, near-IR data of similar quality has been lacking until recently. The new implementation of robotic observations and queue-based operations of other telescopes is changing this situation rapidly.  Suganuma et al.\\ (2006) presented high quality optical and near-IR light curves for 4 Seyfert 1 galaxies obtained over 3 years of intense monitoring. The analysis of these observations, together with some other data available in the literature, showed that the lags observed at near-IR wavelengths are in very good agreement with the location of a putative dusty torus whose inner face should be found at the dust sublimation radius, i.e., at a distance proportional to $L^{1/2}$, where $L$ corresponds to the luminosity of the central source. This behaviour is expected for a torus that intercepts a significant fraction of the near-UV and optical flux from the central source and reradiates it in the near and mid-IR. In this scenario, any flux variations in the central source will be followed by variations at these near-IR wavelengths, with a delay that will be characterised by the location of the torus inner face, which corresponds to the distance at which dust no longer sublimates under the strong glare of the central radiation.  Near-IR emission, however, can be emitted by the colder regions of an accretion disc itself, although there is no clear theoretical predictions of the real extent of the outer regions of the disc. Kishimoto et al.\\ (2008) have presented evidence for near-IR disc emission in Quasars in polarised flux. They find that the spectral energy distribution is very close to $f_{\\nu} \\propto \\nu^{1/3}$, as expected for a Shakura-Sunyaev $\\alpha$ disc. More recently Landt et al.\\ (2011) have shown that emission from an accretion disc and hot dust are required to explain the continuum around 1 micron in a sample of AGN.  In this paper we will present near-IR observations for NGC~3783 and MR~2251-178, and argue that we have detected near-IR variability from their accretion discs. This paper is organised as follows: In Section 2 we characterise our targets; in Section 3 we describe the acquisition and analysis of the data; in Section 4 we give estimates for the host contribution to our nuclear photometry; in Section 5 we present our results in the form of light curves, cross correlation analysis, flux-flux plots and spectral energy distributions; finally, Section 6 discusses our findings and the summary is presented in Section 7. In this article we adopt a concordance $\\Lambda$CDM cosmology with $H_0$ = 70 km/s/Mpc. ",
        "conclusions": "We have obtained high quality near-IR light curves for two nearby AGN: NGC~3783 and MR~2251-178, which are characterised by very different values of their black hole mass. The light curves track variations on time-scales from a few days up to three years for NGC~3783 and four years for MR~2251-178. From the analysis of these data, together with the results already presented in Ar\\'evalo et al.\\ (2008, 2009) we find the following:  \\begin{enumerate}  \\item The Near-IR light curves for both our sources, NGC~3783 and MR~2251-178, show a strong correlation with the optical light curves. \\item For NGC~3783 the optical and near-IR cross correlation analysis suggests that two emitting regions are present in the J-band emission, one consistent with a short lag and another found much further out. We identify these regions as the accretion disc and a dusty torus. The H and K lags ($\\sim 70$ days) are consistent with the presence of dust in a torus located at the sublimation radius. \\item For MR~2251-178 the lags measured between all near-IR and optical wavelengths are consistent with lags of $\\sim 0$ days. This suggests that the variable near-IR emission arises from the accretion disc in this source. \\item From the cross correlation analysis we find no direct evidence for the presence of a dusty torus in MR~2251-178 as expected from the delays predicted by the Suganuma et al.\\ (2006) relation. We checked that the emission is not due to the presence of beaming emission from a radio jet. \\item Flux-flux plots for both sources are consistent with the findings from the cross correlation results. The linear relations observed between the optical and near-IR bands in MR~2251-178 confirm that the variable near-IR emission originates in the accretion disc. However, some near-IR excess seen at low optical fluxes might suggest the presence of a dusty torus. \\item We determined a satisfactory SED representation for NGC~3783 adopting a model of an $\\alpha$-disc illuminated by a central X-ray source plus a dusty torus at a distance consistent with the dust sublimation radius. The model is able to reproduce the optical and near-IR spectral energy distribution as well as other timing constraints. The modelling of MR~2271-178 is not as adequate and requires the presence of a dusty torus. \\item We tentatively interpret the differences in the origin of the near-IR emission in our sources, as well as those reported by Suganuma et al.\\ (2006) and Kishimoto et al.\\ (2009), as the result of the location of the colder regions in the accretion disc with respect to the central illuminating X-ray sources, as originally suggested by Uttley et al.\\ (2003).  \\end{enumerate}"
    },
    "1107/1107.5958_arXiv.txt": {
        "abstract": "We present results from the second part of our analysis of the extended mid-infrared (MIR) emission of the Great Observatories All-Sky LIRG Survey (GOALS) sample based on $5-14\\,\\micron$ low-resolution spectra obtained with the Infrared Spectrograph on \\textit{Spitzer}. We calculate the fraction of extended emission as a function of wavelength for all galaxies in the sample, FEE$_{\\lambda}$, defined as the fraction of the emission that originates outside of the unresolved central component of a source, and spatially separate the MIR spectrum of a galaxy into its nuclear and extended components.  We find that the \\NeII\\ emission line is as compact as the hot dust MIR continuum, while the polycyclic aromatic hydrocarbon (PAH) emission is more extended. In addition, the 6.2 and \\PAHc\\ emission is more compact than that of the \\PAHa, which is consistent with the formers being enhanced in a more ionized medium. The presence of an AGN or a powerful nuclear starburst increases the compactness and the luminosity surface density of the hot dust MIR continuum, but has a negligible effect on the spatial extent of the PAH emission on kpc-scales. Furthermore, it appears that both processes, AGN and/or nuclear starburst, are indistinguishable in terms of how they modify the integrated PAH-to-continuum ratio of the FEE in (U)LIRGs. Globally, the $5-14\\,\\micron$ spectra of the extended emission component are homogeneous for all galaxies in the GOALS sample. This suggests that, independently of the spatial distribution of the various MIR features, the physical properties of star formation occurring at distances farther than 1.5\\,kpc from the nuclei of (U)LIRGs are very similar, resembling local star-forming galaxies with $\\LIR\\,<\\,10^{11}\\,\\Lsun$, as well as star formation-dominated ULIRGs at $z\\,\\sim\\,2$. In contrast, the MIR spectra of the nuclear component of local ULIRGs \\textit{and} LIRGs are very diverse. These results imply that the observed variety of the integrated MIR properties of local (U)LIRGs arise, on average, only from the processes that are taking place in their cores. ",
        "introduction": "\\label{s:introduction}   The mid-infrared (MIR) spectra of galaxies provide key information that allow us to characterize the global properties of their star formation, as well as to develop continuum and emission line diagnostics in order to infer whether they harbor an active galactic nucleus (AGN), with less confusion due to dust extinction compared to observations at optical wavelengths (\\citealt{Lutz1998b}; \\citealt{Genzel1998}; \\citealt{Charmandaris2004}; \\citealt{Brandl2006}; \\citealt{Smith2007}; \\citealt{Armus2007}; \\citealt{Veilleux2009}; \\citealt{Wu2009}; \\citealt{Sales2010}; \\citealt{Petric2011}). However, when studying the integrated emission of galaxies, the spatial distributions of the sources contributing to their observed infrared (IR) spectral energy distributions, such as \\HII\\ regions, photo-dissociation regions (PDRs), cold molecular clouds, and/or non-thermal emission sources, are averaged together. That is, all the intrinsic properties of the different phases of the inter-stellar medium (ISM) are mixed, despite the fact that they can be substantially different (\\citealt{Laurent2000}; \\citealt{Smith2004}). Therefore, even a rough decomposition between the nuclear and disk components of a galaxy can yield valuable information on where the IR luminosity originates, as well as details on the processes associated to the different emission components.  The access to this spatial information is even more vital in the study of luminous and ultra-luminous infrared galaxies, (U)LIRGs, for which it is known that the bulk of their energy production is generated within their central regions and emitted in the IR on scales of a few kpc. This class of galaxies, although not very numerous in the nearby Universe \\citep{Sanders1996}, are responsible for the bulk of the obscured star formation at $z\\,\\geq\\,1$ (\\citealt{LeFloch2005}; \\citealt{PG2005}; \\citealt{Caputi2007}; \\citealt{Magnelli2011}; \\citealt{Murphy2011}). Thus, the study of the spectra of local (U)LIRGs, where the angular resolution achieved by space-born observatories such as \\textit{Spitzer} is sufficient to spatially-resolve their components, is essential for interpreting the internal mechanisms that govern their IR emission and for helping to understand the properties of high redshift (U)LIRGs.  Several MIR spectroscopic studies have already been carried out using ground-based telescopes which provide, for the closest and brightest sources, spatial resolution of a few tens of pc (\\citealt{Soifer2002}; \\citealt{Soifer2003}; \\citealt{DS2010a}). These works have demonstrated that the MIR continuum emission, ionization lines, and polycyclic aromatic hydrocarbon (PAH) features arise from different regions in (U)LIRGs, and that the sources of emission are arranged in more compact configurations than in galaxies with lower IR luminosities. This is consistent with the merger-induced nature of most local (U)LIRGs (\\citealt{Veilleux2002}; \\citealt{Armus2009}) and the funneling of large quantities of gas and dust mass towards their nuclei during the interaction. In addition, MIR spectral maps using the IRS instrument on-board \\textit{Spitzer} have been obtained for 15 local LIRGs revealing an increase of the electron density in their nuclei (\\citealt{PS2010}).  Despite the small number of galaxies studied in these works, some general conclusions about the properties of (U)LIRGs as a whole could be derived. Yet, these trends are confounded by systematics caused by intrinsic variations of processes that dominate at sub-kpc scales. Studying these internal processes is therefore essential if we wish to completely understand the global trends. The Great Observatories All-sky LIRG Survey (GOALS; \\citealt{Armus2009}), which consist of 291 galaxies (202 systems), is the ideal local LIRG sample for this task because it is large enough to allow us to break it up into sub-samples based, for example, on their interaction stage, or the AGN contribution to their MIR emission (\\citealt{Petric2011}). Each sub-sample has enough galaxies in order to perform a robust statistical analysis and examine how the various MIR properties vary from one sub-sample to another and within a given sub-sample. Such an analysis was performed by \\cite{DS2010b}, hereafter Paper~I, who analyzed the compactness of the MIR continuum emission at 13.2$\\,\\micron$ of (U)LIRGs in the GOALS sample and discussed its correlation to several physical properties of the systems such as their \\LIR, MIR AGN fraction, and far infrared (FIR) colors.  The wealth of information available in the $5-14\\,\\micron$ spectral range allow us to further expand our analysis to several other key MIR features. In this second paper we explore how the compactness of the emission due to polycyclic aromatic hydrocarbon (PAH) and atomic lines compares to that of the warm dust continuum and investigate the origin of the observed differences. An analysis regarding the nuclear PAH emission ratios among the GOALS sample will be further addressed in a following study by Stierwalt et al. (2011, in prep.). The paper is structured as follows: in Section 2 we briefly remind the reader of the datasets used and the methodology of our analysis which was described in detail in Paper~I. Our results are reported in Section 3, where we present average MIR spectra of the extended and nuclear emission of the three main types of sources we found in Paper~I. We also explore the role of dust extinction as well as the presence of an active galactic nucleus (AGN) in the observed spatial profiles of the various MIR spectral features. Finally, in Section 4 we present our conclusions.    \\begin{figure*} \\epsscale{1.1} \\plotone{./f1_label.eps} \\caption{\\footnotesize The FEE$_\\lambda$ function, smoothed with a 4-pixel box to reduce noise, are plotted in blue for 3 galaxies that serve as examples of the main types identified in the sample. The Spitzer/IRS $\\sim$5-14$\\mu$m spectrum of each galaxy, scaled to arbitrary units, is also plotted as a pink dashed line for reference. Left panel: NGC~3110 with constant/featureless FEE$_\\lambda$. Middle panel: NGC~1365 displaying a PAH and line extended emission. Right panel: MCG+08-11-002 with silicate-extended emission. }\\label{f:feetype} \\end{figure*} ",
        "conclusions": "\\label{s:summary}   We analyzed the spatial profiles of low spectral resolution $5-14\\,\\micron$ \\textit{Spitzer}/IRS spectra of the GOALS galaxy sample and quantified the spatial extent of their MIR emission, FEE$_\\lambda$. Our work indicates that:  \\begin{itemize}  \\item The \\NeII\\ emission is as compact as the MIR continuum, but the PAHs are more extended in many galaxies. This is in agreement with studies showing that the \\NeII\\ and MIR continuum emissions trace the ionizing stars (i.e., clumpy, current star formation), while the PAH emission is more representative of older, colder, and diffuse stellar populations.  \\item The \\PAHa\\ emission is more extended than that of the 6.2 and \\PAHcs, which is consistent with the formers being enhanced in a more ionized medium while the \\PAHa\\ emission is similar for neutral and ionized molecules.   \\item On average, the MIR continuum becomes more compact than the PAH emission as the MIR is increasingly dominated by an AGN. That is, the presence of an AGN modifies the spatial extent of the hot dust continuum but has a negligible effect on the overall, kpc-scale distribution of the PAH emission. None of the AGN-dominated (U)LIRGs shows a constant/featureless FEE$_\\lambda$ type, implying that the AGN leaves a ``footprint'' in the FEE$_\\lambda$ function of a galaxy. Therefore, the MIR emission of the galaxies showing a constant FEE$_\\lambda$ will not be dominated by an AGN or, in general, show compact MIR continuum emission.  \\item The PAH-to-continuum FEE ratio of (U)LIRGs increases as the global dust temperature, traced by their \\textit{IRAS} log($f_{60\\,\\mu m}/f_{100\\,\\mu m}$) FIR color, becomes warmer. This can be attributed to the presence of a compact, powerful (circum-)nuclear starburst in these galaxies which would increase the temperature of the dust, and also the ratio between the volume of the shells of the photo-dissociation regions --from where PAH emission arises-- and the volume of the hot dust emitting region. ``Normal'' galaxies do not show enhanced PAH-to-continuum FEE ratios but a rather constant value. On the other hand, in (U)LIRGs with large PAH-to-continuum FEE ratios the MIR continuum is more compact than the PAH emission, which is still as extended as in the ``normal'' population since it arises mostly from a diffuse component (their disks). Nonetheless, in terms of the relative spatial distribution of the PAH and continuum emissions, this effect is indistinguishable from that caused by an AGN. Moreover, both processes, \\textit{compact} star formation and AGN activity, coexist in 27\\% of our (U)LIRGs sample.  \\item The intensity of the spectral features and continuum emission of the extended emission component of all (U)LIRGs are very similar, indicating that the properties of the star formation in the external parts of galaxies (disks; $d\\,\\gtrsim\\,1.5\\,$kpc) are uniform. Moreover, they are similar to local starburst galaxies as well as to $z\\,\\sim\\,2$ ULIRGs (selected by their \\textit{Spitzer}/IRAC and MIPS colors), but different from sub-millimeter-selected high-redshift galaxies. On the other hand, the MIR spectra of the unresolved nuclear component of galaxies vary widely.  \\item These results imply that, on average, the diversity seen in the properties of global/integrated MIR spectra of local ULIRGs as well as LIRGs arise only from the processes that are taking place in their cores. The extended star formation appears unaffected by the nuclear activity and it is similar to the more quiescent mode of star formation found in lower IR luminosity systems. We speculate that the outer disks/regions of (U)LIRGs, traced by their PAH spatial profiles, are responsible for most of the cold FIR emission as seen in galaxies with lower IR luminosities, while the nuclei of the most compact (U)LIRGs, as stated in \\cite{DS2010b}, would dominate the MIR continuum and contribute to show overall warmer FIR colors.\\\\     \\end{itemize}"
    },
    "1107/1107.2922_arXiv.txt": {
        "abstract": "We present an analysis of the neutral hydrogen (HI) properties of a fully cosmological hydrodynamical dwarf galaxy, run with varying simulation parameters. As reported by \\citet{Gov10}, the high resolution, high star formation density threshold version of this galaxy is the first simulation to result in the successful reproduction of a (dwarf) spiral galaxy without any associated stellar bulge. We have set out to compare in detail the HI distribution and kinematics of this simulated bulgeless disk with what is observed in a sample of nearby dwarfs.  To do so, we extracted the radial gas density profiles, velocity dispersion (\\eg velocity ellipsoid, turbulence), and the power spectrum of structure within the cold interstellar medium from the simulations.  The highest resolution dwarf, when using a high density star formation threshold comparable to densities of giant molecular clouds, possesses bulk characteristics consistent with those observed in nature, though the cold gas is not as radially extended as that observed in nearby dwarfs, resulting in somewhat excessive surface densities. The lines-of-sight velocity dispersion radial profiles have values that are in good agreement with observed dwarf galaxies, but due to the fact that only the streaming velocities of particles are tracked, a correction to include the thermal velocities can lead to profiles that are quite flat. The ISM power spectra of the simulations appear to possess more power on smaller spatial scales than that of the SMC.  We conclude that unavoidable limitations remain due to the unresolved physics of star formation and feedback within pc-scale molecular clouds. ",
        "introduction": "A traditional problem plaguing the simulation of disk galaxies \\citep[e.g.][and references therein]{Thacker01,som03, Abad03a, Gov04, Gov07, Rob04, Bai05,Okam05, San09, Stinson10}, within a cosmological context, has been the inability to recover successfully the properties of a truly ``late-type'' disk and, in particular, those with essentially no associated stellar bulge, similar to classical galaxies such as M33.  Recent work by \\citet{Gov10}, though, has produced what appears to be exactly such a bulgeless dwarf, via the imposition of a higher density threshold for star formation (100~cm$^{-3}$, as opposed to 0.1~cm$^{-3}$, as adopted in the aforementioned earlier generations of simulations), and mass resolution that allows one to identify individual star forming regions.\\footnote{The higher star formation density threshold can only be applied because the high resolution of the simulation, coupled with heating from the UV background, ensures fragmentation does not occur at unresolved scales.} The primary dwarf in their analysis\\footnote{In addition to the supplementary re-simulation of DG1 described in \\S~2.} (DG1) forms a shallow central dark matter profile and possesses a pure exponential stellar disk of radial scale $r_d$$\\sim$1~kpc, with a stellar bulge-to-disk ratio $B/D$$\\approx$0.04 as determined from the $i$ band light profile.  In what follows, we extend this work and examine in detail the cold neutral hydrogen (HI) gas content of the simulated dwarf DG1 along with its low star formation threshold analog, DG1LT, and an updated version of DG1 (called, nDG1) which employs high-temperature metal-line cooling and enhanced supernova energy feedback to compensate for the additional cooling. Our goal is simple: to determine if their HI gas properties agree with recent observational data to an equally successful degree as the stellar component. Studies such as the The HI Nearby Galaxy Survey \\citep[THINGS][]{wal08} provide excellent high resolution (spectral and spatial) data against which to compare our simulations. The gas properties of the simulations are compared directly with several of the most recent relevant empirical datasets \\citep{Stan99,Tam09, Obr10}, in order to assess both their strengths and weaknesses.  The cold gas in galaxies is linked directly to underlying star formation processes and associated interstellar medium (ISM) physics; any successful model of galaxy formation should adopt a holistic approach, examining both the gas and star properties in consort.  We describe the basic properties of our simulations, before detailing the analyses undertaken; we will present results pertaining to the radial distribution of cold gas within the disks associated with DG1, DG1LT, and nDG1, spatially-resolved velocity dispersion maps of the cold gas, and the spatial distribution of power encoded within the structure of the ISM. We end with a summary of our findings, discussing both the strengths and weaknesses of the current simulations. ",
        "conclusions": "One immediate concern arising from our analysis relates to the issue of extracting ``neutral hydrogen'' from the simulations' ``cold gas'' (which in some sense consists of both molecular and neutral hydrogen).  Because the high-density regions within the simulation have densities more akin to molecular, rather than neutral, clouds, it is important to explore the definition of ``neutral'' employed here.\\footnote{In large part, this was motivated by the fact that in ``column density space'', these high-density regions possess column densities close to 10$^{22}$~cm$^{-2}$, higher than those observed in nature; this is a limitation of the conversion employed within {\\textsc{Gasoline}}.} To do this, we re-generated HI moment maps, but now restricting the gas included to only those particles with densities near the classical value of $\\sim$0.1~cm$^{-3}$.  As expected, this eliminated the unrealistically high neutral hydrogen column densities in the highest density regions, but at the expense of leading to vertical density profiles that bore little resemblance to the Gaussian profiles observed in nature \\citep{Obr10}.  Such an extreme ``cut'' to the definition of neutral hydrogen also led to a radial profile that bore little resemblance to an exponential.  We found no density cut which impacted favourably on the observable properties of DG1. For these simulations, because density and temperature are closely correlated in the relevant regime (T$\\simlt$30000~K; $\\rho$$\\simgt$0.001~cm$^{-3}$), the above analysis is degenerate to cuts in volume density or temperature.  It is important to note that the primary process responsible for driving bulk properties in the simulation is the star formation and feedback prescription.  \\citet{Gov10} demonstrated that star formation had a larger effect on the rotation curve of our simulated galaxy than resolution (see their Figure 5). The gas properties presented in this paper are primarily the result of the star formation prescription, and thus it is imperative to use a star formation and feedback prescription that is physically motivated.  Until metal-dependent H$_2$ creation and cooling is added to the simulations, it is not clear how much HI, as opposed to H$_2$, should be present in the simulation, how it might be distributed as a function of radius, and what impact it will have on the resulting disk.   After applying  the physically-motivated $\\sim$8~km/s thermal broadening to the Cartesian coordinates of the SPH particles' streaming motion, the inferred \\it characteristic \\rm velocity dispersions for the cold gas were a reasonable match to those observed in nature (albeit, at the unavoidable expense of recovering any correlation between velocity dispersion and galactocentric radius and/or global star formation in the disk, in addition to the imposition of an isotropic velocity ellipsoid to the cold gas, and the young stars which form from this gas).  Beyond the aforementioned issue of the lack of a self-consistent treatment of molecular cooling processings on sub-parsec scales, one must also be aware that at the resolutions of these simulations, we are still missing unresolved star forming regions and associated turbulence.  The nature of these missing sources is an active area of debate, but magnetorotational instability (MRI) is one of the favoured mechanisms capable of providing a non-negligible amount of turbulence \\citep[e.g.][]{WangAbel09, PionOstr07, MacLow09}  Enhancing the supernovae energy feedback, as was done for simulation nDG1, at these resolutions, had a marginal impact on the SPH particles' streaming velocities (at the $\\sim$20\\% level), which in turn meant that its impact on the velocity dispersion profiles was also minimal. This is not surprising, as the increased energy deposition was used in order to offset the effect of the newly included high-temperature metal-line cooling. Without the inclusion of extra SN energy, the additional cooling that comes from metal lines leads to more star formation than in the case of DG1.  As shown by \\citet{Oh2011}, the stellar mass of DG1 is in good agreement with galaxies at similar halo masses, as observed by THINGS.  If high-temperature metal-line cooling had been added with $\\epsilon$SN held constant, nDG1 would have overproduced stars for galaxies in a comparable halo mass range. However, the enhanced feedback seems to have steepened the spatial power spectrum of the cold ISM of nDG1 relative to DG1, making it less consistent with the power spectrum observed for the SMC. It is unclear, however, how the power spectrum varies with the instantaneous SFR and if this result holds across time.  Capturing all the relevant ISM physics necessary to recover the full spectrum of turbulence sources at pc and sub-pc scales remains an outstanding challenge.  Despite these limitations, the simulated dwarf galaxies presented here have been shown to possess bulk characteristics consistent with those observed in nature, including adherence to scaling relations such as the size-luminosity, size-velocity, and luminosity-velocity \\citep{brooks2011}. Additionally, the star formation and feedback prescription used in these simulations has been shown to result in a realistic mass-metallicity relationship as a function of time, and consume gas at a rate that reproduces the incidence rate and metallicities of both QSO-Damped Lyman Alpha (DLA) and GRB-DLA systems \\citep{Brooks2007, Pontzen2008, Pontzen2009}.  Hence, it is clear that our simulations remain extremely successful in recovering many of the {\\it global} optical and dynamical properties of realistic bulgeless dwarfs.  That is, although the microphysics of the ISM cannot be fully captured at the force resolutions that must be used currently in cosmological simulations, this does not largely impact the bulk macrophysics such as the rotation curves (stellar and dark matter mass profiles), angular momentum content, etc.  On the other hand, we have seen that higher resolutions and adoption of more realistic physics for star formation leads to simulated galaxies that better reproduce the properties of observed galaxies \\citep[e.g.,][]{Booth2007, RobertsonKravtsov2008, TaskerBryan2008, Saitoh2008, CeverinoKlypin2009, Gov10}.  The work presented here highlights paths for future improvement in the implementation of ISM physics in cosmological simulations, and provides useful tests for reassessment once metal-dependent H$_2$ cooling and star formation has been added to {\\sc Gasoline} and other cosmological simulation codes."
    },
    "1107/1107.2113_arXiv.txt": {
        "abstract": "Recent discoveries of strongly misaligned transiting exoplanets pose a challenge to the established planet formation theory which assumes planetary systems to form and evolve in isolation. However, the fact that the majority of stars actually do form in star clusters raises the question how isolated forming planetary systems really are. Besides radiative and tidal forces the presence of dense gas aggregates in star-forming regions are potential sources for perturbations to protoplanetary discs or systems. Here we show that subsequent capture of gas from large extended accretion envelopes onto a passing star with a typical circumstellar disc can tilt the disc plane to retrograde orientation, naturally explaining the formation of strongly inclined planetary systems. Furthermore, the inner disc regions may become denser, and thus more prone to speedy coagulation and planet formation. Pre-existing planetary systems are compressed by gas inflows leading to a natural occurrence of close-in misaligned hot Jupiters and short-period eccentric planets. The likelihood of such events mainly depends on the gas content of the cluster and is thus expected to be highest in the youngest star clusters. ",
        "introduction": "\\label{sec:intro} The discoveries of retrograde or strongly misaligned transiting exoplanets like WASP-17b and others \\citep{Heetal08,Gillon09,Joetal09,Naritaetal09,Pontetal09,Pontetal10,Winetal09a,Winetal09b,wasp17b_01,Triaudetal2010} pose a challenge for established planet formation theories. These classically assume the birth of planetary systems out of collapsing protostellar cloud fragments. While contracting from sub-parsec to milliparsec (or hundreds of AU) scale the conservation of its initial angular momentum causes the cloud to spin and flatten, while part of its mass settles towards the center, eventually igniting the hydrogen burning in the stellar core. Given such a protostar with a circumstellar disc that contains a certain amount of dust, solid particles collide and stick together, forming larger particles. This coagulation leads to bodies (cores) of sufficient mass to accrete surrounding material (dust and even gas) by gravity until no more disc material is available. This mechanism is probably the dominant process of the formation of the Solar System, including the Earth \\citep{2008ASPC..398..235M}. As an alternative, gravitational instability as a forming mechanism for giant planets and brown dwarfs is being discussed \\citep{2004ApJ...610..456B}, and at least for brown dwarfs it has actually been shown to work \\citep{StHuWi07,Thiesetal2010}. More recent variations of these models assume gravitational instability of the solid phase only or dust trapping in vortices as a speed-up for coagulation.  Being isolated from any external perturbation, everything inside this protostar-disc system spins in the same direction, thus the forming planets orbit their host star the same way. This simple model has been severely questioned by the findings mentioned above. The transiting exoplanet WASP-17b \\citep{wasp17b_01}, for example, has been found to have a sky-projected inclination of the orbital plane normal against the stellar spin axis of $148.5\\pma{5.1}{4.2}$ degrees. A number of other mis-aligned transiting exoplanets have been discovered \\citep{Heetal08,Gillon09,Joetal09,Naritaetal09,Pontetal09,Pontetal10,Winetal09a,Winetal09b,wasp17b_01}. These discoveries suggest that misalignment between planetary orbits and the spin of their host star (hereafter called spin-orbit misalignment) are quite common, at least among close-in transiting planets for which the spin-orbit alignment can be measured by using the Rossiter-McLaughlin Effect \\citep{Rossiter1924,McLaughlin1924}. The nearby $\\upsilon$~And planetary system exhibits even mutually inclined planetary orbits \\citep{2010ApJ...715.1203M}.  Such planetary orbits cannot be explained by the standard planet formation model according to which the conservation of angular momentum forces any coplanar system to remain coplanar forever. There have been recent suggestions that may shed some light into the possible mechanism. \\cite{FabTre07,Naozetal2011} deduce a formation scenario for misaligned ``hot jupiters'' through long-term mutual perturbations of two planets (or an inner planet and an outer brown dwarf; for simplicity, we use the term ``planet'' for all substellar companions in this paper) orbiting the same star. Given an initial mutual inclination between 40\\grad\\ and 140\\grad\\ the Kozai-Lidov mechanism leads to secular oscillations of eccentricity vs. inclination. Tidal friction circularises the orbit of the inner planet if its periastron falls below a few stellar radii, eventually leaving the planet on a close-in orbit with near-random inclination with respect to the stellar spin. This scenario, however, requires a sufficient initial mutual inclination of both planets.  Another possible mechanism is tilting the spin axis of the star. Direct stellar encounters within a stellar radius (about 0.01~AU) can be effectively ruled out, because they are improbable, even within stellar clusters \\citep{TKT05, TK07} and would destroy the disc. Secular transfer of angular momentum between the protostellar spin and the protoplanetary disc via magnetic fields have been investigated by \\citet{Laietal11}.  Multi-stage or episodic accretion of circumstellar material may provide another viable mechanism for misaligned planetary systems. The underlying idea is that accretion from different sources (i.e. gas filaments, accretion envelopes) from different directions may sever the classically assumed spin-orbit correlation. The star itself may get its angular momentum from a different accretion event than the bulk of circumstellar material. \\citet{BLP10} have analysed turbulent accretion events using data from \\citet{BBB03}. They deduce a high frequency of misaligned planets as a consequence of multi-directional accretion. However, due to the limited simulation data available, they were forced to make simplified assumptions to what happens within the close vicinity of a star. Our work aims to take a deeper insight into such scenarios by self-consistent calculations of close interactions of a circumstellar disc with encountered new material, and by also treating the effects on pre-existing planets.  As \\citet{PTS08,OPE10} have shown, encounters between stars and pre-/protostellar objects do occur frequently in dense stellar environments and may lead to significant accretion bursts. Although the precise mechanisms proposed \\citep{PTS08} differ from ours, their calculations emphasise the importance of encounter events to accretion processes.  In this paper the effects of gas capture onto a pre-existing circumstellar disc are studied. In Section \\ref{sec:model} we depict the physical scenario while Section \\ref{sec:methods} briefly describes the numerical methods used in our computations. The results are presented in Section \\ref{sec:results}. ",
        "conclusions": "\\label{sec:conclusions} We have analysed the general effects of inclined gas capture onto a pre-existing circumstellar disc upon passage through a dense gaseous reservoir. The secondary gas capture does not destroy a pre-existing disc even if the new material inflows at a highly inclined angle. Rather, the inner regions of the disc (inside about 30 AU) become denser, while some of the outer disc material is scattered away. The major result is the formation of an inclined combined circumstellar disc from captured material and the pre-existing disc. After capturing a few 0.01~\\tmsun\\ into an initially inclined annulus both structures tend to align with each other within a few 10~000 years after the encounter. An annulus with an initial inclination of $135\\grad$ aligns to about $35\\grad$ within 20~000 years. Due to conservation of angular momentum this results in a net shift of more than $90\\grad$ in our model, and even more for lower-mass pre-existing discs or disc-less stars. If planets form from the resulting inclined disc one naturally arrives at misaligned planets.  The resulting disc may have two chemically distinct radial regions, the inner region which is associated with the star, and the outer region composed from the captured material.  In addition, the basic effect of capture onto a pre-existing planetary system have been estimated. Orbits typically shrink and become strongly eccentric in response to the contact with captured gas and subsequent planet-planet interaction, while their orbital planes can be dramatically altered thus naturally leading to (even mutually) misaligned short-period planets with eccentric orbits.  Important constraints are given by recent observations of apparent inclinations between circumstellar discs and their host stars. In a most recent analysis \\citet{Watsonetal2011} have found no significant misalignment between the normal vectors of debris discs and the spin axis of low-mass to solar-type stars. In particular, the projected inclinations between the star and its disc typically differ by 5--10 degrees, and 15--45 degrees in a few cases. Their results do not necessarily contradict our scenario since even moderate disc tilts and warps may lead to a Kozai resonance in a planetary system born out of such a disc. Following the mechanism described by \\citet{FabTre07} and \\citet{Naozetal2011} near-random misalignment of a close-in planet and the spin of its host star may result.  Future studies will also consider protoplanetary discs with embedded planets and systems with a brown dwarf on a wide orbit (about 100--200~AU). In addition, the consequences for dust content and mixing of differently composed material will be analysed.  The mechanism described in this paper provides a natural scenario for the formation of misaligned planetary systems in gas-rich dense star-forming regions. While other models have been proposed our scenario requires rather simple assumptions, and may even provide the initially misaligned planetary orbits required by the recent model of \\citet{Naozetal2011}."
    },
    "1107/1107.1326_arXiv.txt": {
        "abstract": "Relativistic jets are ubiquitous in astrophysical systems that contain compact objects. They transport large amounts of energy to large distances from the source, and their interaction with the ambient medium has a crucial effect on the evolution of the system. The propagation of the jet is characterized by the formation of a shocked \"head\" at the front of the jet which dissipates the jet's energy and a cocoon that surrounds the jet and potentially collimates it. We present here a self consistent, analytic model that follows the evolution of the jet and its cocoon, and describes their interaction. We show that the critical parameter that determines the properties of the jet-cocoon system is the dimensionless ratio between the jet's energy density and the rest-mass energy density of the ambient medium. This parameter, together with the jet's injection angle, also determines whether the jet is collimated by the cocoon or not. The model is applicable to relativistic, unmagnetized, jets on all scales and may be used to determine the conditions in AGNs jets as well as in GRBs or microquasars. It shows that AGN and microquasar jets are hydrodynamically collimated due to the interaction with the ambient medium, while GRB jets can be collimated only inside a star and become uncollimated once they breakout. ",
        "introduction": "Relativistic jets are observed in many astrophysical systems which host compact objects, such as radio galaxies, gamma-ray bursts (GRBs) and microquasars. These jets appear in a variety of lengths, durations and energy scales, and propagate in different types of media. However, despite the differences in their characteristics, which are many orders of magnitudes apart, the interactions of the jets with their surroundings lead to similar results, such as energy injection into the ambient medium and its feedback on the jet. Understanding these common phenomena can provide valuable insights on puzzles such as what type of medium can collimate the jet and how much energy is injected into the surrounding matter. These in turn can be used to study the heating of the interstellar and intergalactic medium by active galactic nuclei (AGNs) and the fate of jets and stellar envelops of collapsing massive stars during gamma-ray bursts.   The interaction of a jet with the external medium was studied extensively in many different scales, using analytic and numerical methods. Initially it was studied in the context of radio loud AGNs \\citep[][]{BlandRees74,Scheuer74}. They showed that the propagation of the jet generates a double bow-shock structure at the head of the jet. Energy and matter that enter this structure are pushed aside due to a high pressure gradient and create a hot cocoon around the jet. The cocoon, in turn, applies pressure on the jet and compresses it. \\citet{BegelCiof89}  calculated the head's velocity and the cocoon expansion rate, assuming that the cocoon and the head are supported by the ram-pressure of the ambient medium. Their calculations assumes Newtonian head velocities and rely on the knowledge of the cross-section of the head, which they estimated phenomenologically using  the size of the radio lobes observed around jets of radio loud AGNs.  \\citet{MW01} calculated the Lorentz factor of a relativistic head assuming a conical jet (i.e., no collimation). \\citet{Matzner03} extended the model of \\citet{BegelCiof89} to include both Newtonian and relativistic head velocities, assuming a conical jet as in \\citet{MW01}. Later, \\citet{LazzBeg05} used this model to calculate the opening angle and the properties of the jet at breakout from a surface of a collapsing star. They accounted for the reduction in the opening angle of the jet due to collimation, but ignored the dissipation of energy in shocks that form within the jet as a result of such a collimation. This led to an incorrect dependence of the opening angle on the cocoon's pressure. An attempt to include the effects of these shocks was made by \\citet{Morsony07}.  Apart from the analytical studies the numerical approach was also given a lot of attention over the years. Simulations were conducted in the context of extra galactic jets \\citep[e.g.][]{Marti95,Marti97,Aloy99,Hughes02} as well as GRB jets that propagate inside a star \\citep[e.g.][]{Zhang03,Morsony07,Mizuta09}. All these works showed the basic features discussed above, i.e. the formation of a jet head and a hot cocoon. A collimation shock at the base of the jet is also evident in the simulations, whenever the jet is collimated by the cocoon. However, despite the extensive numerical and analytical efforts, to this date there is no simple analytic description of the time evolution of the jet-cocoon system. As a result, there is no quantitative understanding of the collimation process, and its effect on the jet angle.  The goal of this work is to provide a simple self consistent, analytic description, for the time evolution of an unmagnetized relativistic jet and the cocoon that forms  around it, when it propagates in a density profile that is suitable for most astrophysical systems. The key point in our analysis is the treatment of the collimation shock that forms at the base of the jet. This shock is an inevitable consequence of collimation in a supersonic jet, as it dissipates part of the jet's energy, generating the pressure needed to counterbalance the cocoon's pressure. It is crucial for the proper modeling of the system, since it sets the width of the jet and controls its propagation velocity. We base our solution on the earlier analysis of \\citet{BegelCiof89,Matzner03}, and incorporate the geometry of the collimation shock as it is analyzed by \\citet{KomisFall97, BL09}. We find that the jet evolution can be of two types: collimated and un-collimated according to the strength of the jet - cocoon interaction and the collimation shock. We quantify the transition between these regimes and derive models for the jet and the cocoon evolution in each regime. We test the validity of our approximations by comparing the results from our model with output of various numerical simulations.  The paper is constructed as follow: In \\S2 we present a general nontechnical description of the system, its ingredients and the different stages in its evolution. In \\S3 we present our model and derive the criterion that distinguishes between the two collimation regimes. We discuss the geometry of the collimation shock in each regime and analyze how it affects the temporal evolution of the system. We also provide approximated analytic solutions in each of the regimes. We compare our analytical model to hydrodynamic simulations in \\S4. In \\S5 we briefly consider jet collimation in various astrophysical environments. Thorough study of specific systems is lengthy and will be presented in following papers. In \\S6 we summarize the main results. The full list of results, that specifies the behavior of the physical quantities in the entire parameter space is summarized in table 1 and in appendix B. ",
        "conclusions": "In this work we present an analytical study of the propagation of a relativistic hydrodynamic jet in an external medium with a general density profile. The interaction of the jet with the medium results in the formation of a shocked \"head\" at the front of the jet, and an over pressured cocoon with a rather uniform distribution of energy surrounding the jet. When the pressure of the cocoon, $P_c$, is larger than the internal pressure of the jet, it compresses the jet and leads to the formation of an oblique shock at the base of the jet. Under some conditions, determined in this work, the shock converges to the jet's axis, and the jet becomes collimated. This changes the jet's properties, and affects the conditions in the cocoon. Our model follows the evolution of the jet, the jet's head, the cocoon and the collimation shock, given three initial parameters: the jet's luminosity, $L_j$, the jet's opening angle at the injection point, $\\theta_0$, and the density of the ambient medium, $\\rho_a(z)$. We determine the condition for collimation, and provide a self-consistent, time dependent, solution to the system's parameters that depends on these three initial parameters alone. The main results of our analysis can be summarized as: \\begin{itemize} \\item The jet's evolution can be classified into two regimes: a collimated and an un-collimated, according to the strength of interaction of the jet with the ambient medium. The two regimes are distinguished by the parameter $\\tilde L$ and by $\\theta_0$ (or $\\theta_j$), which represents the ratio of the jet's energy density  to the rest-mass energy density of the ambient medium at the location of the head. \\item When $\\tilde L<\\theta_0^{-4/3}$ the interaction is strong and the jet is collimated. In this regime the collimation shock converges on the jet's axis at some point, $\\hat z$,  below the head. Above $\\hat z$ the jet is cylindrical to a good approximation and its width is estimated as $\\hat z\\theta_0/2$. This width implies that the  Lorentz factor of the collimated ejecta satiafies $\\Gamma_{j1}=\\theta_0^{-1}$. \\item If $\\tilde L\\ll1$, the head of the collimated jet is non-relativistic ($\\beta_h\\ll1$). In this limit the cocoon's expansion velocity is proportional to the head's velocity, and it has a constant opening angle $\\theta_c\\simeq\\theta_0$. \\item When $\\tilde L\\gg\\theta_0^{-4/3}$ the collimation shock fails to converge and it remains at the edge of the jet. The jet in this regime in un-collimated and it remains conical to a good approximation. It has a coaxial structure of a fast unshocked spine surrounded by a denser layer of shocked material with lower Lorentz factor. \\item When $\\tilde L\\gg\\theta_0^{-4}$, $\\Gamma_h>\\theta_0^{-1}$, and the head is not causal in the transverse direction. Here only a fraction $\\sim (\\Gamma_h\\theta_0)^{-1}$ of the jet energy flows into the cocoon. This energy is enough to produce a mildly relativistic pressure in the cocoon, and accelerate cocoon's edge to a velocity $\\beta_c\\sim1$. \\item When $\\tilde L>4\\Gamma_j^{-4}$ the reverse shock becomes Newtonian and the head no longer affects the jet's propagation. Matter continue to stream into the cocoon only from the forward shock, and it generates a relativistic pressure in the cocoon. \\item The value of $\\tilde L$ can change with time, even for a constant $L_j$ and $\\theta_0$, according to the slope of the density profile. In a density profile $\\rho_a\\sim z^{-\\alpha}$ with $\\alpha>2$, $\\tilde L$ increases with time. This corresponds to an acceleration of the jet's head. In such cases a collimated jet opens up and become un-collimated once $\\tilde L$ becomes larger than $\\theta_0^{-4/3}$. The opposite evolution take place when $\\alpha<2$. In the case of $\\alpha=2$, $\\tilde L$ is constant, the head maintains its velocity, and the opening angle of a collimated jet remains constant. \\item The choice of the initial injection radius of the jet, $z_*$, can affect the jet's breakout time, where larger values of $z_*$ increase this time. This effect becomes important when the ratio of $z_*$ to the stellar radius $\\gtrsim10^{-3}$. Numerical simulations typically use values of $z_*$ which are above this limit and therefore their resultant breakout times are affected by this choice. \\end{itemize}   Our model provides a general frame to study the properties of relativistic jets in different media. It can be used to examine various phenomena, such as the minimal energy required by the jet to break out of a boundary surface at a finite distance, like the edge of a star. It can also be used to test the energy feedback into the stellar envelope in the case of a GRB jet that penetrates a star, or the IGM in the case of AGN jets. This may help understands better issues such as the problem of IGM heating. Our model confirms that GRB jets, which form inside a star, are collimated before they break out and become uncollimated afterwards. Microquasars and X-ray binary jets, however, may start uncollimated, but the interaction with the ISM leads to their collimation beyond $\\sim2\\cdot10^{-3}$ pc. We also show that quasar jets are collimated as well at large distances from their sources ($\\gtrsim10$ kpc). We stress that in our calculations we assume that the magnetic fields in the jet, in the region where the jet undergoes the collimation and above, are dynamically unimportant and therefore can be ignored. High magnetization in this region will alter our results, for example by preventing the formation of the collimation shock, but will keep the other basic properties of the model, i.e. the formation of the cocoon and the collimation of the jet.   We are in dept to A. Mizuta for providing us results from his numerical simulations. We also thank Y. Lubarsky for useful discussions, and to the anonymous referee for helpful comments. O.\\,B.\\ and T.\\ P. were supported by the Israel Center for Excellence for High Energy Astrophysics and by an ERC advanced research grant. E.\\,N.\\ was supported in part by the Israel Science Foundation (grant No.\\ 174/08) and by an EU International Reintegration Grant. R.\\, S. \\ is partially supported by IRG \\& ERC grant, and a Packard fellowship.   \\renewcommand{\\theequation}{A-\\arabic{equation}} \\setcounter{equation}{0}  %"
    },
    "1107/1107.4029_arXiv.txt": {
        "abstract": "Known bright S stars, recognized as such by their enhanced s-process abundances and C/O ratio, are typically members of the asymptotic giant branch (AGB) or the red giant branch (RGB).  Few modern digital spectra for these objects have been published, from which intermediate resolution spectral indices and classifications could be derived. For published S stars we find accurate positions using the Two-Micron All Sky Survey (2MASS), and use the FAST spectrograph of the Tillinghast reflector on Mt. Hopkins to obtain the spectra of 57 objects.  We make available a digital S star spectral atlas consisting of 14 spectra of S stars with diverse spectral features.  We define and derive basic spectral indices that can help distinguish S stars from late-type (M) giants and carbon stars.   We convolve all our spectra with the Sloan Digital Sky Survey (SDSS) bandpasses, and employ the resulting $gri$ magnitudes together with 2MASS $JHK_s$ mags to investigate S star colors. These objects have colors similar to carbon and M stars, and are therefore difficult to distinguish by color alone.  Using near and mid-infrared colors from IRAS and AKARI, we identify some of the stars as intrinsic (AGB) or extrinsic (with abundances enhanced by past mass-transfer).  We also use $V$ band and 2MASS magnitudes to calculate a temperature index for stars in the sample.    We analyze the proper motions and parallaxes of our sample stars to determine upper and lower limit absolute magnitudes and distances, and confirm that most are probably giants. ",
        "introduction": "\\label{sec:introduction}  When a low to intermediate mass star exhausts its core supply of hydrogen, contraction and heating of the core causes the outer layers of the star to expand and cool, whereby the star becomes a red giant (RG) with an average effective temperature of 5000\\,K or lower.  Once the core temperature rises to about 3$\\times 10^8$\\,K, helium burning begins, and the stellar surface contracts and heats creating the horizontal branch (for Population II) or red clump (for Population I) stars on the color-magnitude diagram.  Completion of core helium burning and the start of helium shell burning returns the star to a luminous, cooler phase as it evolves onto the asymptotic giant branch (early- or E-AGB) where the star greatly expands and the hydrogen shell is (almost) extinguished.  Later, as the He-burning shell approaches the H-He discontinuity, a phase of double shell burning begins.  Recurrent thermal instabilities in this thermally-pulsing asymptotic giant branch (TP-AGB) phase are driven by helium shell flashes, periodic thermonuclear runaway events in the He-shell \\citep{1983ARA&A..21..271I} The energy released by these pulses expand the star, whereby the hydrogen shell is again basically extinguished during some time. A third dredge up might then occur, and afterward the hydrogen shell will re-ignite (and a new thermal pulse will occur).  During these short-lived pulses, nucleosynthesis products from combined H-shell and He-shell burning are dredged up to the outer layers of the star.  The complex details of the dredge-up and its outcome are strongly mass-dependent (\\citealt{2005ARA&A..43..435H} and references therein), but generally speaking, surface abundance enhancements result in C, He, and the $s$-process elements, including Ba, La, Zr, and Y.  S and Carbon (C) stars are traditionally thought to be members of the TP-AGB.  In addition to possibly elevated C/O ratios, these stars exhibit other unique spectral characteristics as a result of the dredge-up.  While M giants often exhibit titanium oxide (TiO) absorption bands, both carbon and zirconium have a higher affinity to free oxygen.  Therefore, a higher C/O ratio also results in the disappearance of TiO bands and the appearance of zirconium oxide (ZrO) bands.  S stars are therefore distinguished primarily by their ZrO bands, while C stars show strong C$_2$ and  CN bands. The S star spectral class is divided into three subtypes (in order of increasing carbon abundance):  MS, S, and SC stars.  Classification into one particular class is difficult, but is generally based on the C/O ratio and the relative strengths of TiO, ZrO, and CN molecular absorption bands.  Overall, a star in the S star class is generally defined as having a C/O ratio (r) of $0.5 < r < 1.0$ and strong ZrO absorption features \\citep{1993A&A...271..463J,2010arXiv1011.2092V}.  Pure S stars are those that show only ZrO bands, and no TiO bands \\citep{1978MNRAS.184..127W}. Classification is complicated by the fact that increased intensity of the ZrO bands may be due to an excess of Zr, rather than being exclusively tied to the C/O ratio \\citep{2003IAUS..210P..A2P}. Furthermore, many AGB S and C giants are highly variable, and since the strength of molecular bands depends on the effective temperature and surface gravity of the star, spectral type can vary significantly with time.  S stars whose spectra show lines of the short-lived element $s$-process technetium (Tc) are known as \"intrinsic\" S stars, and are most likely be members of the TP-AGB. $^{99}$Tc is only produced by the $s$-process, while the star is on the AGB.  Since the half-life of $^{99}$Tc is only 2.13 x 10$^5$ years, and the average duration of the AGB phase of stellar evolution is roughly 1 Myr, stars exhibiting significant abundances of technetium are almost certainly members of the AGB \\citep{2000A&A...360..196V}. By contrast, technetium-poor S stars are known as \"extrinsic\" S stars. These stars are likely to be part of binary systems and their unusual chemical abundances probably originated in a past mass-transfer episode \\citep{2000A&A...360..196V}.  The present extrinsic S star likely once accreted $s$-process rich material from its companion, a TP-AGB star at the time that has since evolved into a white dwarf. Extrinsic S stars therefore display enriched $s$-process elements, despite having never produced these elements themselves.  Most known extrinsic S stars are assumed to be members of the red giant branch (RGB).  They can sometimes be distinguished from intrinsic S stars on the basis of color, because they do not necessarily show the red excess characteristic of stars on the AGB.  While all currently known extrinsic S stars are thought to be giants, these definitions open up the possibility of extrinsic S stars on the main sequence (hereafter referred to as dwarf S stars) that have previously accreted $s$-process material and carbon from a companion but whose Tc has since decayed. Analogous to the dwarf C stars \\citep{Green91}, a faint S star could be shown to be a dwarf if it has a sufficiently large parallax or proper motion; a large proper motion predicts a tangential velocity greater than the Galactic escape speed unless the object is faint and nearby.  S stars are very similar to C stars, with the major difference being a slightly lower C/O ratio.  Faint high galactic latitude carbon (FHLC) stars were once assumed be giant stars at large distances.  In 1977, however, G77-61 was discovered to have high proper motion and an upper limit absolute magnitude of +9.6, therefore making it the first known carbon dwarf star \\citep{1977ApJ...216..757D}.  In this case, the term dwarf refers to a star on the main sequence.  \\citet{Green91} discovered four other dwarf carbon (dC) stars using techniques that selected known carbon stars for high proper motion.  Now, well over 100 dwarf carbon stars are known \\citep{2004AJ....127.2838D}. The local space density of dCs is far higher than all types of C giants combined \\citep{1992ApJ...400..659G}.  Therefore, contrary to previous assumptions, dwarf C stars are the numerically dominant type of carbon star in the Galaxy.  Since dC stars are now known to be so common, we have begun a search for dwarf S stars.  There is little reason to believe that dwarf S stars should not exist.  In fact, they may be more abundant than dC stars, since less carbon enhancement is required to produce an S dwarf as compared to a C dwarf.  On the other hand, if the range of abundance ratios that produces the characteristic features of S stars is narrow, they may be quite rare.  Furthermore, the range of C/O ratios that produces S star spectral features may be different for giants and dwarfs, given the higher gravities and larger temperature ranges in dwarfs.  We use medium-resolution digital spectra of known S stars (which are probably giants) to generate corresponding Sloan Digital Sky Survey (SDSS) colors, to see whether their colors might be sufficiently distinctive to search for additional S stars in the SDSS with reasonable efficiency.  We also use the FAST spectra to investigate molecular band spectral indices for classification.  Our paper is organized as follows: in \\S\\ref{sec:obs} we briefly review the observations and our data reduction and processing procedures.  In section \\S\\ref{sec:spatlas} we describe the production of a spectral atlas of S stars, and derive spectral indices from the molecular band strengths in \\S\\ref{sec:specind}. In \\S\\ref{sec:colors} we describe the generation of SDSS colors and analyze the overall colors of the FAST sample in order to determine whether an efficient color selection for S stars exist.  In \\S\\ref{sec:intext} we discuss techniques for distinguishing intrinsic and extrinsic S stars and apply them to the FAST sample.  In \\S\\ref{sec:motion} we analyze the parallaxes and proper motions of S stars in the sample to characterize upper and lower limit magnitudes and distances.  We use colors to generate approximate temperature indices for S stars in \\S\\ref{sec:temp}.  In \\S\\ref{sec:sum} we summarize our results and present our conclusions. ",
        "conclusions": "\\label{sec:sum}  Using a sample of 57 medium-resolution S star spectra taken with the FAST spectrograph, we create a spectral atlas of S stars comprised of 14 objects that span a range of spectral types within the MS, S, and SC classes.  The atlas is published as a collection of 1-D FITS files via this journal.  After generating $g$ and $r$ SDSS magnitudes from the spectra, we find that the SDSS magnitudes in the $g$ and $r$ bands, combined with $J$, $H$, and $K_s$ magnitudes from the 2MASS catalog for S stars may with further confirmation allow for reasonably efficient color selection of these objects from the SDSS and 2MASS catalogs.    We use previously published data to identify some of the stars in the sample as intrinsic or extrinsic stars, and find that the fraction of extrinsic S stars in the sample is approximately $54\\%$.  We also assign temperature indices to the stars in the sample based on the M star scale of temperature indices.  We find that much of the S star sample falls at temperature index 4 or above, meaning that the effective temperatures of most of these stars are well below 5000 K. We also analyze objects in the FAST sample with detectable parallaxes and proper motions to generate absolute magnitude limits, as well as lower limit distances based on assumed dwarf magnitudes.  This analysis bears out our initial assumption that  the FAST sample is primarily composed of giant stars, either on the AGB or RGB."
    },
    "1107/1107.4154.txt": {
        "abstract": "We consider the effect of the coupled variations of fundamental constants on the nucleon magnetic moment. The nucleon $g$-factor enters into the interpretation of the measurements of variations in the fine-structure constant, $\\alpha$,  in both the laboratory (through atomic clock measurements) and in astrophysical systems (e.g. through measurements of the 21 cm transitions). A null result can be translated into a limit on the variation of a set of fundamental constants, that is usually reduced to $\\alpha$. However, in specific models, particularly unification models, changes in $\\alpha$ are always accompanied by corresponding changes in other fundamental quantities such as the QCD scale, $\\Lambda_{\\scriptscriptstyle\\rm QCD}$. This work tracks the changes in the nucleon $g$-factors induced from changes in $\\Lambda_{\\scriptscriptstyle\\rm QCD}$ and the light quark masses.  In principle, these coupled variations can improve the bounds on the variation of $\\alpha$ by an order of magnitude from {\\em existing} atomic clock and astrophysical measurements. Unfortunately, the calculation of  the dependence of $g$-factors on fundamental parameters is notoriously model-dependent. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Any definitive measurement of a temporal or spatial variation in a fundamental constant, such as the fine-structure constant $\\alpha$, would signal physics beyond the standard model, and in particular a violation of the equivalence principle which is one of the foundations of general relativity. In many cases, such an observation would indicate the existence of a new light (usually scalar) degree of freedom \\cite{Bek}. Indeed, there has been considerable excitement during the last decade over the possible time variations in $\\alpha$ from observations of quasar absorption systems \\cite{webb,murphy3,chand,quast,murphy07,chand3}.  In effectively all unification models of non-gravitational interactions, and certainly in models in which one imposes gauge coupling unification at some high energy scale, a variation in $\\alpha$ is invariably accompanied by variations in other gauge couplings \\cite{ekow,co}. In particular, variations in the strong gauge coupling, $\\alpha_s$, will induce variations in the  QCD scale, $\\Lambda_{\\scriptscriptstyle\\rm QCD}$, as can be seen from the low energy expression for $\\Lambda_{\\scriptscriptstyle\\rm QCD}$ when mass thresholds are included \\begin{equation} \\Lambda_{\\scriptscriptstyle\\rm QCD} = \\mu \\left(\\frac{m_{\\rm c} \\, m_{\\rm b} \\, m_{\\rm t}}{\\mu^3} \\right)^{\\frac{2}{27}} \\, \\exp\\left[-\\frac{2\\pi}{9\\alpha_s(\\mu)} \\right] \\,, \\end{equation} % for a renormalization scale $\\mu > m_{\\rm t}$ up to the unification scale  \\cite{ekow,co,lang}, where $m_{\\rm c,b,t}$ are the masses of the charm, bottom, and top quarks. Because fermion masses are proportional to $hv$ where $h$ is a Yukawa coupling and $v$ is the Higgs vacuum expectation value (vev), variations in Yukawa couplings will also affect variations in $\\Lambda_{\\scriptscriptstyle\\rm QCD}$ so that \\begin{eqnarray}\\label{DeltaLambda} \\frac{\\Delta \\Lambda_{\\scriptscriptstyle\\rm QCD}}{\\Lambda_{\\scriptscriptstyle\\rm QCD}} &= &R \\, \\frac{\\Delta \\alpha}{\\alpha} + \\frac{2}{27} \\left(3 \\, \\frac{\\Delta v}{v} + \\frac{\\Delta h_{\\rm c}}{h_{\\rm c}} + \\frac{\\Delta h_{\\rm b}}{h_{\\rm b}} +  \\frac{\\Delta h_{\\rm t}}{h_{\\rm t}} \\right) \\,. \\end{eqnarray} Typical values for $R$ are of order 30 in many grand unified theories, but there is considerable model-dependence in this coefficient \\cite{dine}.  Furthermore, in theories in which the electroweak scale is derived by dimensional transmutation, changes in the Yukawa couplings (particularly the top Yukawa) lead to exponentially large changes in the Higgs vev. In such theories, the Higgs expectation value is related to the Planck mass, $M_{\\rm P}$, by~\\cite{co} \\beq\\label{rad1} v\\sim M_{\\rm P} \\exp\\left(- \\frac{2 \\pi c}{ \\alpha_{\\rm t}}\\right) , \\eeq where $c$ is a constant of order 1, and $\\alpha_{\\rm t} = h_{\\rm t}^2/4\\pi$. For $c \\sim h_{\\rm t} \\sim 1$, \\beq\\label{enhance2} {\\Delta v \\over v} \\sim S {\\Delta h_{\\rm t} \\over h_{\\rm t}} , \\eeq with $S \\sim 160$, though there is considerable model-dependence in this value as well. For example, in supersymmetric models, $S$ can be related to the sensitivity of the $Z$ gauge boson mass to the top Yukawa, and may take values anywhere from about 80 to 500 \\cite{eos}. This dependence gets translated into a variation in all low energy particle masses \\cite{ds}.  In addition, in many string theories, all gauge and Yukawa couplings are determined by the expectation value of a dilaton and we might expect \\cite{co} \\beq\\label{h-alpha} {\\Delta h \\over h} = {1 \\over 2} {\\Delta \\alpha \\over \\alpha} , \\eeq assuming that all Yukawa couplings vary similarly, so that they all reduce to $h$. Therefore, once we allow $\\alpha$ to vary, virtually all masses and couplings are expected to vary as well, typically much more strongly than the variation induced by the Coulomb interaction alone.   Irrespective of the purported observations of a time variation in $\\alpha$, many experiments and analyses have led to limits on possible variations \\cite{uzan,uzan2}. Furthermore, the use of coupled variations has led to significantly improved constraints in a wide range of environments ranging from big bang nucleosynthesis \\cite{co,ichikawa,wett,cnouv,df,landau,grant}, the Oklo reactor \\cite{opqccv,wett2}, meteoritic data \\cite{opqccv,opqmvcc,wett2}, the microwave background \\cite{cmb,landau} and stellar evolution \\cite{ek}.  This article explores the possibility that the strongest existing limits on the fine-structure constant, namely those derived from atomic clock measurements, can also be enhanced by considering such coupled variations. We expect the effect of induced variations in $ \\Lambda_{\\scriptscriptstyle\\rm QCD}$ and the light quark masses to enter through the nucleon magnetic moment. Existing experimental limits on $\\alpha$ from atomic clock experiments assume constant $\\mu_{\\rm p,n}$. Indeed, limits on the variations of quark masses in units of the QCD scale, i.e. $m_{\\rm q}/\\Lambda_{\\scriptscriptstyle\\rm QCD}$, from atomic clock measurements have been derived \\cite{flam,lattice result ref 0402098}. Given a (model-dependent) calculation of the nucleon magnetic moment (or equivalently its $g$-factor), we can derive sharper bounds on the variation of $\\alpha$ from existing data. Unfortunately, because of the model-dependence, we find that while the limits are generally improved (by as much as an order of magnitude), there is considerable uncertainty in the precise numerical limit. As a corollary, we apply our results to astrophysical measurements such as those which rely on the 21 cm line which also depends on $\\mu_{\\rm p,n}$.  The article is organized as follows: In section 2, we outline the procedure of obtaining limits on $\\alpha$ from atomic clock experiments. In particular, we examine the detailed dependence on the nuclear $g$-factors which will be subject to variation. In section 3, we derive the dependence of the nucleon magnetic moment on $\\Lambda_{\\scriptscriptstyle\\rm QCD}$ and the light quark masses. Because there is no unique (or rigorous) method for calculating baryon magnetic moments, we consider several different approaches.  The most straightforward employs the constituent quark model.  Surprisingly, this model is quite effective in matching the observed baryon magnetic moments. Even within this broad approach, our result will depend on the calculation of the nucleon mass, as well as the calculation of the constituent quark mass; each carrying a significant degree of uncertainty.  We also consider an approach based on chiral perturbation theory, and a method based partially on lattice results.  In section 4, we apply these results to atomic clock measurements and derive ``improved\" limits on the variation of $\\alpha$.  Finally, in section 5, we extend these results to measurements involving the 21 cm line and summarize our results.      %%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\subsection{Astrophysical systems}  Several different types of observations of astrophysical systems involving quasar absorption spectra are subject to a similar analysis that has been applied to atomic clocks. Indeed, there are four distinct combinations of physical parameters which depend on $g_{\\rm p}$.  \\begin{itemize} \\item The comparison of UV heavy element transitions with the hyperfine H\\,{\\sc i} transition allows one to set constraints on \\begin{equation} x\\equiv\\alpha^2 g_{\\mathrm{p}} \\mu, \\end{equation} since the  optical transitions are simply proportional to $R_\\infty$. It follows that constraints on the time variation of $x$ can be obtained from high resolution 21~cm spectra compared to UV lines, e.g., of Si\\,{\\sc ii}, Fe\\,{\\sc ii} and/or Mg\\,{\\sc ii}. The recent detection of 21~cm and molecular hydrogen absorption lines in the same damped Lyman-$\\alpha$ system at $z_{\\mathrm{abs}}=3.174$ towards SDSS J1337+3152 constrains~\\cite{x21cm} the variation $x$ to \\begin{equation} \\Delta x/x=-(1.7\\pm1.7)\\times10^{-6},\\qquad z=3.174. \\end{equation}  \\item The comparison of the  H\\,{\\sc i} 21~cm hyperfine transition to the rotational transition frequencies of diatomic molecules allows one to set a constraint on \\begin{equation} y \\equiv g_{\\mathrm{p}}\\alpha^2 \\end{equation} The most recent constraint~\\cite{y-murph} relies on the comparison of two absorption systems determined both from H\\,{\\sc i} and molecular absorption. The first is a system at $z=0.6847$ in the direction of TXS~0218+357 for which the spectra of CO(1-2), \\super{13}CO(1-2), C\\super{18}O(1-2), CO(2-3), HCO\\super{+}(1-2) and HCN(1-2) are available. They concluded that \\begin{equation} \\Delta y/y=(-0.16\\pm0.54)\\times10^{-5}, \\qquad z=0.6847. \\end{equation} The second system is an absorption system in the direction of PKS~1413+135 for which the molecular lines of CO(1-2), HCO\\super{+}(1-2) and HCO\\super{+}(2-3) have been detected. The analysis led to \\begin{equation} \\Delta y/y=(-0.2\\pm0.44)\\times10^{-5},\\qquad z=0.247. \\end{equation}  \\item The ground state, ${}^2\\Pi_{3/2} J=3/2$, of OH is split into two levels by $\\Lambda$-doubling and each of these doubled levels is further split into two hyperfine-structure states. Thus, it has two ``main'' lines ($\\Delta F=0$) and two ``satellite'' lines ($\\Delta F=1$). Since these four lines arise from two different physical processes ($\\Lambda$-doubling and hyperfine splitting), they enjoy the same Rydberg dependence but different $g_{\\mathrm{p}}$ and $\\alpha$ dependences. By comparing the four transitions to the H\\,{\\sc i} hyperfine line, one can set a constraint on \\begin{equation} F\\equiv g_{\\mathrm{p}}(\\alpha^2/\\mu)^{1.57}. \\end{equation} Using the four 18~cm OH lines from the gravitational lens at $z\\sim0.765$ toward PMN~J0134-0931 and comparing the H\\,{\\sc i} 21~cm and OH absorption redshifts of the different components allowed one to set the constraint~\\cite{OH-3} \\begin{equation} \\Delta F/F=(-0.44\\pm0.36\\pm1.0_{\\text{syst}})\\times10^{-5},\\qquad z=0.765, \\end{equation} where the second error is due to velocity offsets between OH and H\\,{\\sc i} assuming a velocity dispersion of 3~km/s. A similar analysis~\\cite{darling} in a system in the direction of PKS~1413+135 gave \\begin{equation} \\Delta F/F=(0.51\\pm1.26)\\times10^{-5},\\qquad z=0.2467. \\end{equation}  \\item The satellite OH~18~cm lines are conjugate so that the two lines have the same shape, but with one line in emission and the other in absorption.  This behavior has recently been discovered at cosmological distances  and it was shown~\\cite{OH-1} that a comparison between the sum and difference of satellite line redshifts probes the variation of \\begin{equation} G \\equiv g_{\\mathrm{p}}(\\alpha^2/\\mu)^{1.85}. \\end{equation} From the analysis of a system at $z\\sim0.247$ towards PKS~1413+135, it was concluded~\\cite{kanekar1} that $|\\Delta G/G| =(2.2\\pm3.8) \\times10^{-5}$, while a newer analysis~\\cite{kanekar1bis} gave \\begin{equation} |\\Delta G/G| =(-1.18\\pm0.46) \\times10^{-5}. \\end{equation} It was also applied to a nearby system~\\cite{kanekar2}, Centaurus~A, to give $ |\\Delta G/G| < 1.16\\times10^{-5}$ at $z\\sim0.0018$. \\end{itemize}  These constraints are summarized in Table~\\ref{tab:astro}.  \\begin{table}[htb!] \\caption{Constraints on the variation of different combinations of $g_{\\rm p}$, $\\mu$ and $\\alpha$ from astrophysical observations.} \\label{tab:astro} \\small \\begin{center} \\begin{tabular}{|l|ccc|cc|} \\hline Combination & $\\lambda_{g_{\\rm p}}$ & $\\lambda_\\mu$ & $\\lambda_\\alpha$ & Constraints (yr$^{-1}$) & redshift  \\\\ \\hline\\hline $x=g_{\\mathrm{p}} \\alpha^2 \\mu$ & $1$ &$1$ & $2$ & $-(1.7\\pm1.7)\\times10^{-6}$ & $3.174$ \\\\ $y=g_{\\mathrm{p}}\\alpha^2$ & $1$ & $0$ & $2$ & $(-0.16\\pm0.54)\\times10^{-5}$ & $0.6847$ \\\\ & & & & $(-0.2\\pm0.44)\\times10^{-5}$ & $0.247$ \\\\ $F=g_{\\mathrm{p}}(\\alpha^2/\\mu)^{1.57}$ & $1$ & $-1.57$ & $3.14$ & $(-0.44\\pm0.36\\pm1.0_{\\text{syst}})\\times10^{-5}$ & $0.765$ \\\\ & & & & $(0.51\\pm1.26)\\times10^{-5}$ & $0.2467$ \\\\ $G=g_{\\mathrm{p}}(\\alpha^2/\\mu)^{1.85}$ & 1 &$-1.85$ & 3.70 & $(-1.18\\pm0.46)\\times10^{-5}$& $0.247$ \\\\ & & & & $(0\\pm1.16)\\times10^{-5}$ & $0.0018$ \\\\ \\hline \\end{tabular} \\end{center} \\normalsize \\end{table}  \\subsection{Astrophysical constraints}  In contrast to our analysis of atomic clocks, we cannot combine the astrophysical observations because they have been obtained from different systems at different redshifts and at different spatial locations. However, as we have done previously (but without the $g_{\\rm n}$ and $b$ terms), we show in Table~\\ref{tab:factorastro} the enhancement factor for the analysis of the 4 types of combinations of absorption spectra. We emphasize that the enhancement factor is always larger than unity (except for $y$ in the EOMS and $\\chi$PT+QCD models).  As last example of the power of coupled variations, Table~\\ref{tab:astrocte} compares the constraints on the variation of $\\alpha$ that can be obtained under the assumption that $g_{\\rm p}$ and $\\mu$ are constant with the assumption of coupled variations based on unification. As one can see, in many cases the limits are improved by an order of magnitude.  \\begin{table}[htb!] \\caption{Value of the parameter $C_\\alpha$ for the 4 combinations of constants that can be constrained by astrophysical observations, assuming $R=30$ and $S=160$ (left) and value of the enhancement factor $C_\\alpha /\\lambda_\\alpha$ (right). } \\label{tab:factorastro} \\small \\begin{center} \\begin{tabular}{|l|cccc|} \\hline & $x$ & $y$ & $F$ &  $G$ \\\\ \\hline\\hline A & $34.10$ & $15.49$ & $-12.60$ & $-17.25$ \\\\ B1  & $20.72$ & $2.10$ & $-25.98$ & $-30.64$ \\\\ B2 & $24.96$  & $6.35$ & $-21.73$ & $-26.39$ \\\\ B3  & $23.52$ & $4.91$ & $-23.18$ & $-27.83$ \\\\ C &  $31.21$  & $12.59$ & $-15.49$ & $-20.14$ \\\\ \\hline HBw/oD  & $2.46$ & $-15.76$ & $-43.23$ & $-47.77$ \\\\ HBwD  &  $27.00$ & $8.78$ & $-18.69$ & $-23.23$ \\\\ EOMS &  $17.63$ & $-0.60$ & $-28.07$ & $-32.61$ \\\\ \\hline $\\chi$PT+QCD & $16.96$ & $-1.26$ & $-28.73$ & $-33.27$ \\\\ \\hline \\end{tabular} \\begin{tabular}{|l|cccc|} \\hline & $x$ & $y$ & $F$ &  $G$ \\\\ \\hline\\hline A   &  $17.05$  & $7.74$ & $-4.01$ & $-4.66$ \\\\ B1  & $10.36$  & $1.05$ & $-8.27$ & $-8.28$ \\\\ B2  & $12.48$  & $3.18$ & $-6.92$ & $-7.13$ \\\\ B3  &  $11.76$ & $2.45$ & $-7.38$ & $-7.52$ \\\\ C   &  $15.60$ & $6.30$ & $-4.93$ & $-5.44$ \\\\ \\hline HBw/oD  & $1.23$ & $-7.88$ & $-13.77$ & $-12.91$ \\\\ HBwD  & $13.50$ & $4.39$ & $-5.95$ & $-6.28$ \\\\ EOMS & $8.81$ & $-0.30$ & $-8.94$ & $-8.81$ \\\\ \\hline $\\chi$PT+QCD & $8.48$ &$-0.63$ & $-9.15$ & $-8.99$ \\\\ \\hline \\end{tabular} \\end{center} \\normalsize \\end{table}  \\begin{table}[htb!] \\caption{Comparison of the constraints obtained from astrophysical systems with and without assumption on unification for model A.} \\label{tab:astrocte} \\small \\begin{center} \\begin{tabular}{|l|ccc|} \\hline Combination & independent (yr$^{-1}$) & correlated (yr$^{-1}$) & redshift  \\\\ \\hline\\hline $x$ & $(-8.50\\pm8.50)\\times10^{-7}$ & $(-4.98\\pm4.98)\\times10^{-8}$ & 3.174 \\\\ $y$ & $(-0.8\\pm2.7)\\times10^{-6}$ & $(-1.03\\pm3.49)\\times10^{-7}$ & $0.6847$ \\\\ & $(-1.0\\pm2.2)\\times10^{-6}$ & $(-1.29\\pm2.84)\\times10^{-7}$ & $0.247$ \\\\ $F$ & $(-1.40\\pm3.38)\\times10^{-6}$ & $(3.49\\pm8.44)\\times10^{-7}$& $0.765$ \\\\ & $(1.62\\pm4.01)\\times10^{-6}$& $(-0.40\\pm1.00)\\times10^{-6}$ & $0.2467$ \\\\ $G$ & $(-3.19\\pm1.24)\\times10^{-6}$ & $(6.84\\pm 2.67)\\times10^{-7}$& $0.247$ \\\\ & $(0\\pm3.14)\\times10^{-6}$ & $(0\\pm6.73)\\times10^{-7}$& $0.0018$ \\\\ \\hline \\end{tabular} \\end{center} \\normalsize \\end{table}   \\subsection{Discussion}  In this article, we have discussed the effect of a correlated variation of fundamental constants, focusing on the gyromagnetic factors $g_{\\rm p}$ and $g_{\\rm n}$. These parameters are particularly important to interpret electromagnetic spectra, and thus to derive constraints on the variation of fundamental constants from atomic clock experiments and from quasar absorption spectra. As discussed, there is an important model-dependence in the computation of the gyromagnetic factors in terms of the quark masses and QCD scale.  When applied to the interpretation of atomic clock experiments, we have shown that in general the constraints on the variation of $\\alpha$ are sharper than that under the assumption that $g_{\\rm p}$, $g_{\\rm n}$, $b$ and $\\mu$ are constant, but this is not a systematic conclusion as we have exhibited models in which the variation of $\\alpha$ stays the same or is even weaker due to cancellations in the sensitivity to $\\alpha$. The constraints on the variation of $\\alpha$ should then be taken with care. In many cases, they may be stronger than reported, but they may be weaker as well.  This points to the need to better understand the fundamental physics needed to calculate baryon magnetic moments. Any limit which depends on $g_{\\rm p,n}$ will be subject to the type of uncertainties discussed here.  Fortunately, the tightest constraint arises from the Hg-Al clock experiments, that does not depend on $g_{\\rm p}$, $g_{\\rm n}$, $b$ or $\\mu$. As a consequence, we have been able to independently set a bound on the variation of $hv$ from the combination of the other experiments. While this bound is still model-dependent, we have shown that it is always smaller than \\begin{equation} \\left|\\frac{(hv)^.}{hv}\\right|<2.0\\times10^{-15}\\,{\\rm yr}^{-1} \\end{equation} for the models we have considered in this article.  Our analysis also applies to astrophysical system and to quasar absorption spectra. We have shown that the enhancement factor is almost always larger than unity.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1107/1107.1867_arXiv.txt": {
        "abstract": "Using \\chandra\\ imaging spectroscopy and Very Large Array (VLA) L-band radio maps, we have identified radio sources at P$_{1.4GHz}$$\\geq$5$\\times$10$^{23}$ \\radiounits\\ and X-ray point sources (XPSs) at L$_{0.3-8keV}$$\\geq$5$\\times$10$^{42}$ \\lum\\ in L$>$L$^*$ galaxies in 12 high-redshift (0.4$<$$z$$<$1.2) clusters of galaxies. The radio galaxies and XPSs in this cluster sample, chosen to be consistent with Coma Cluster progenitors at these redshifts, are compared to those found at low-$z$ analyzed in \\citet{hart09}. Within a projected radius of 1 Mpc of the cluster cores, we find 17 cluster radio galaxies (11 with secure redshifts, including one luminous FR~II radio source at $z$$=$0.826, and 6 more with host galaxy colors similar to cluster ellipticals). The radio luminosity function (RLF) of the cluster radio galaxies as a fraction of the cluster red sequence (CRS) galaxies reveals significant evolution of this population from high-$z$ to low-$z$, with higher power radio galaxies situated in lower temperature clusters at earlier epochs. Additionally, there is some evidence that cluster radio galaxies become more centrally concentrated than CRS galaxies with cosmic time. Within this same projected radius, we identify 7 spectroscopically-confirmed cluster XPSs, all with CRS host galaxy colors. Consistent with the results from \\citet{2009ApJ...701...66M}, we estimate a minimum X-ray active fraction of 1.4$\\pm$0.8\\% for CRS galaxies in high-$z$ clusters, corresponding to an approximate 10-fold increase from 0.15$\\pm$0.15\\% at low-$z$. Although complete redshift information is lacking for several XPSs in $z$$>$0.4 cluster fields, the increased numbers and luminosities of the CRS radio galaxies and XPSs suggest a substantial (9-10 fold) increase in the heat injected into high redshift clusters by AGN compared to the present epoch. ",
        "introduction": "\\label{sec:p2_intro}  Clusters of galaxies are special environments in which to study the evolution of galaxies and active galactic nuclei (AGN) because these massive structures contain a hot, diffuse intracluster medium (ICM) that is detectable in the X-ray.  The disturbed appearance of extended radio emission from individual cluster galaxies \\citep[e.g., NGC 1265 in the Perseus cluster;][]{1986ApJ...301..841O} was the first indication that an ICM existed \\citep[e.g.,][]{1972Natur.237..269M} and early X-ray observations detected it at T$=$10$^{7-8}$ K \\citep[][and references therein]{1986RvMP...58....1S}.  With the high-resolution of the \\chandra\\ Advanced CCD Imaging Spectrometer (ACIS), it was found that the radio-lobes of nearby AGN in the brightest cluster galaxy (BCG) appear to be coincident with cavities in X-ray surface brightness \\citep[e.g., Perseus cluster;][]{2002MNRAS.331..369F}. The physical connection suggested by this spatial coincidence allows an indirect measurement of the kinetic power produced by these AGN.  The connection between AGN and galaxy evolution is not well-understood, but theoretical results suggest that the absence of extremely massive galaxies is due to the truncation of star formation \\citep[e.g.,][]{2004ApJ...608..752B}.  AGN have been implicated as the mechanism to ``shut-off\" star-formation by heating the surrounding gas and/or ejecting the gas reservoirs necessary for galaxy growth during outburst episodes \\citep[e.g.,][]{2006MNRAS.370..645B}. Heating by AGN has also been suggested to account for anomalies in the scaling relationship between cluster X-ray luminosity and temperature \\citep[][and references therein]{2004MNRAS.348.1078B}. Particularly in low-mass clusters (i.e., clusters with low X-ray temperatures), an ``entropy\" floor is observed; \\citet{1999Natur.397..135P} suggest that preheating of the ICM is required to explain this feature.  This entropy floor may be due to the heat injection by AGN in these clusters over time.  In general the X-ray surface brightness of a cluster ICM is well-fit by a $\\beta$-model, which is similar to a King profile.  However, some clusters display a strong peak of X-ray emission at their centers.  Assuming that thermal bremsstrahlung is the dominant cooling process in cluster cores, the cooling time for the X-ray emitting gas is shorter than the Hubble time in these ``cool-core clusters\".  Therefore, prodigious amounts of star formation ($\\geq$100~M$_{\\sun}$/year) are expected; however, spectroscopic X-ray observations fail to detect line emission from this gas at kT$<$1~keV \\citep{2003ApJ...590..207P}, nor do optical observations reveal significant amounts of star formation \\citep[e.g.,][]{1989AJ.....98.2018M,1993AJ....105..417M}.  The lack of strong evidence for wide-spread star formation in cluster cores suggests that a source of heat is preventing the ICM from cooling in cluster cores.  AGN are likely contributors to ``feedback\" processes that are required to prevent this large-scale cooling.  The connection between AGN and their surrounding environments, both small and large, proves to be a difficult problem to simulate. Many cosmological simulations are large volume boxes in which computational efficiency is balanced by the need to resolve small structures, such as galaxies, and to track cluster scale-properties that are several orders of magnitudes larger than galaxies.  Thus, in practice, some cosmological simulations inject a pre-determined amount of ``heat\" at the centers of evolving clusters to simulate pre-heating of the cluster cores.  Hydrodynamical simulations of the plasma-filled lobes created by the AGN in cluster BCGs can recreate the observed X-ray structures \\citep[e.g.,][]{2004ApJ...611..158R} and halt the formation of strong cool-core clusters; however, these simulations in general cannot efficiently distribute the heat throughout the cluster core \\citep[e.g.,][]{2007ApJ...671..171V}. One potential solution to this difficulty is that more than just the one AGN in the BCG contribute to heating the ICM.  Previous studies have found statistical excesses of radio galaxies \\citep[e.g.,][]{2007ApJS..170...71L} and X-ray point sources, or XPSs, in the vicinity of galaxy clusters \\citep[e.g.,][]{2005ApJ...623L..81R, 2005AA...430...39C}. The higher surface density of AGN in and around clusters brings up some interesting questions: (1) Is the enhanced surface density simply a direct result of a higher density of galaxies? (2) Are cluster AGN triggered into activity as they pass through the dense cluster cores? (3) Is the cluster environment especially suited in some specific way to create AGN activity?  Radio galaxies and XPSs are clearly associated with AGN activity at the core of their host galaxies.  In order to consistently study cluster AGN, observations obtained in at least three wavelength regimes are required to detect these objects and to understand the energy feedback processes between them and the surrounding ICM.  In \\citet{hart09}, hereafter Paper I, we examined the cluster radio galaxy and XPS population in eleven \\lowz\\ clusters.  Within 1 Mpc from the cluster center, we detected 20 radio galaxies at \\rlim\\ and 8 XPSs at \\lxbb$\\geq$\\xlim, all spectroscopically-confirmed as cluster members with host galaxies at L$\\geq$L$^*$. These cluster sources are hosted by cluster ``red sequence\" (CRS) ellipticals, but are more centrally concentrated than the CRS galaxies as a whole. The cluster XPSs in our low-$z$ cluster sample show little evidence for obscuration based both on their optical colors, which are consistent with the CRS, and also on the absence of significant obscuration in the X-ray.  Based on the optical and X-ray properties of these cluster XPSs, we suggested that these sources are most similar to low-luminosity, high-energy peaked BL Lacertae (BL Lac) objects. Based on the extrapolations of the radio luminosity function (RLF) of radio galaxies and the X-ray luminosity function (XLF) of BL Lacs to lower radio and X-ray luminosities, we estimate that almost all CRS galaxies are radio-loud at P$_{1.4GHz}$$\\geq$10$^{21.4}$ \\radiounits\\ and X-ray-loud at \\lxbb$\\geq$10$^{40}$ \\lum. Therefore, most CRS galaxies possess AGN jets which can deposit heat into the ICM as they move through the inner cluster regions.  Paper I concluded that the radio-loud CRS population between 10$^{21.4}$$<$P$_{1.4GHz}$$<$10$^{23.5}$ \\radiounits\\ can inject $\\sim$55\\% of the total energy to the ICM if we assume a shallow scaling law of AGN jet power to radio luminosity, namely P$_{jet} \\propto$ L$_r^{0.5}$ from \\citet{2008ApJ...686..859B}.  More recently \\citet{2010ApJ...720.1066C} extended the \\citet{2008ApJ...686..859B} sample to lower-power radio galaxies and measured a steeper slope for this scaling relationship (P$_{jet} \\propto$ L$_r^{0.7}$), which decreases the estimated heat input by low-luminosity AGN to 30\\%.  In Paper I the number of confirmed cluster radio galaxies is much larger than cluster XPSs and so cluster radio galaxies will dominate the ICM heating regardless of the interpretation of XPSs.  In this paper, we continue our study of cluster AGN begun in Paper I with the goal of studying the evolution of radio galaxies and XPSs in galaxy clusters. We examine twelve \\comboz\\ clusters to detect radio galaxies at P$_{1.4GHz}$$\\geq$5$\\times$10$^{23}$ \\radiounits\\ and XPSs at L$_{0.3-8.0keV}$$\\geq$5$\\times$10$^{42}$ \\ergs\\ within a 1 Mpc projected radius of the cluster cores to determine the distribution of AGN luminosities, their locations relative to the ICM, and their numbers and luminosities as a function of redshift. In \\S\\ref{sec:p2_rtc} we describe the cosmological simulations that guide the selection of our cluster sample, uniquely designed to permit an accurate evolutionary study.  In \\S\\ref{sec:p2_obs}, we detail the X-ray, radio, and optical datasets used to measure the ICM X-ray temperatures, to identify potential cluster XPSs and radio sources, and to obtain optical colors of typical cluster galaxies and AGN host galaxies.  In \\S\\ref{sec:p2_obs_results} we present the sample of potential cluster radio galaxies and XPSs, summarize their optical properties through composite color-magnitude diagrams, and describe the typical colors of cluster radio galaxies and XPSs relative to the CRS.  In \\S\\ref{sec:p2_rg}--\\ref{sec:p2_rlf} we present the RLF of cluster radio galaxies out to $z$$=$1.1 and compare it to other cluster RLFs.  In \\S\\ref{sec:p2_xps} we examine the XPS population in our cluster fields out to $z$$=$1.2 to determine if any statistical XPS excess above the expected background is observed and estimate the X-ray active fraction of CRS galaxies. In \\S\\ref{sec:p2_implications} we discuss our results and in \\S\\ref{sec:p2_conclusions} we summarize our conclusions from this study.  X-ray luminosities are defined for the rest-frame (0.3--8.0 keV) bandpass, unless otherwise noted. Radio powers are defined for a 1.4 GHz rest-frame frequency.  Throughout this paper, we use H$_0$ = 70~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.70$, and $\\Omega_{M}=0.3$. ",
        "conclusions": "\\label{sec:p2_conclusions}  As a companion paper to \\citet[][Paper I]{hart09}, we have presented an extensive investigation of the multiwavelength nature of AGN in a sample of 22 clusters of galaxies at 0.2$<$$z$$<$1.2, specifically chosen to be most similar to Coma Cluster progenitors at those redshifts.  Using the X-ray temperature of the ICM as a proxy for cluster mass, we chose massive clusters (or high X-ray temperature clusters) at lower redshifts and less massive clusters (or lower X-ray temperature clusters) at higher redshifts to investigate the X-ray point source (XPS) and radio galaxy populations. In both wavelength regimes, the limiting luminosities were chosen to sample only the AGN populations in L$\\geq$L$^*$ galaxies.  Within 1 Mpc of the cluster center, we detect 19 radio galaxies at \\plim\\ in our \\comboz\\ clusters (11 of which are spectroscopically-confirmed cluster members and 8 of which are likely cluster members because they have colors similar to the cluster red sequence (CRS). Because 8 CRS AGN do not have confirming spectroscopic redshifts, we conservatively reduce the total number of radio galaxies in this sample by 2 to 17 based upon the foreground/background AGN contamination we found within the bounds of the CRS in Paper I. Similarly, we detect 7 spectroscopically-confirmed cluster XPSs.  Both the X-ray and radio AGN have host galaxies with colors consistent with CRS galaxies. There is some evidence that the radio galaxies become more centrally concentrated in these clusters at lower redshift. Among the radio galaxies we have found a clear FR~II morphology source at $z$$=$0.8 with P$_{1.4GHz}$$=$7.7$\\times$10$^{26}$ \\radiounits, while no FR~IIs were detected at P$_{1.4GHz}$$>$10$^{25}$~\\radiounits\\ in our low-$z$ sample.  In addition to this one luminous FR-II, we detect several other radio galaxies at P$_{1.4GHz}$$>$10$^{25}$ \\radiounits\\ in our $z$$>$0.4 clusters, including several BCGs.  These sources are all more powerful than any we found at low-$z$.  While previous work \\citep[e.g.,][]{1991ApJ...371...49E,1991ApJ...367....1H,2002AJ....124.1239H} has discovered clusters of galaxies around clear FR~II morphology and power level sources at $z$$>$0.4, here we started with rich clusters at those redshifts and have discovered an FR~II. Therefore, there is no doubt that the morphology and power level of some cluster radio galaxies is changing between $z$$=$1 and 0. Specifically, this study provides further support for a phenomenological model in which the power levels of subsequent outbursts of radio galaxies in clusters diminishes rapidly between these redshifts \\citep{1991ApJ...371...49E,2002AJ....124.1239H}, the same range of redshifts over which the ICM density and temperature dramatically increase (see Figure 1). Evidence presented herein and in Paper I that radio galaxies are more centrally concentrated in their clusters than the CRS elliptical galaxy population as a whole is further evidence for a link between AGN activity and ICM density \\citep[see also][]{1981ApJ...245..375S}. It is possible that this AGN/ICM relationship may provide the necessary negative feedback mechanism to prevent runaway cooling in cluster cores \\citep[e.g.,][]{2008MNRAS.386.1309M}.  A comparison between the RLF of our \\lowz\\ radio sources versus those in our \\comboz\\ sample reveals a larger number of high-power sources at earlier epochs, evidence for positive evolution of these sources in Coma Cluster progenitors.  A comparison between the RLF of our \\lowz\\ cluster radio sources to local (z$<$0.09) Abell cluster radio sources shows that the two populations are similar, suggesting little evolution between \\lowz\\ and $z$$=$0. The clear density and luminosity evolution of radio galaxies found in our sample of cluster AGN is quite different from the absence of evolution found by \\citet{1999AJ....117.1967S} using a flux-limited sample of EMSS clusters.  We interpret this lack of agreement as support for our novel, physically-based method for selecting clusters for study, the ``Road to Coma'' concept presented in the \\S\\ref{sec:p2_rtc}. We suggest that both AGN evolution results are correct and that \\citet{1999AJ....117.1967S} failed to find cluster AGN evolution because their high-$z$ sample has very similar ICM properties to the low-$z$ sample of \\citet{1996AJ....112....9L}, to which they compared AGN populations. For example, the hot, $z$$=$0.8 cluster, MS 1045-03 (see Figure~\\ref{fig:cluster_sim}) is a member of the \\citet{1999AJ....117.1967S} high-$z$ sample and has a similar radio galaxy population to many $z$$=$0 rich Abell clusters. Our sample was selected to avoid that pitfall and we conclude that sample selection is as critical for cluster galaxy and AGN studies as it is for other evolutionary investigations.  We can be far less definitive concerning the evolution of cluster X-ray AGN due to the large number of XPSs in our high-$z$ cluster fields and to the absence of redshift information for most of them. However, there are sufficient redshifts (7) for CRS XPSs to show that the passive, absorption-line AGN population that was identified in Paper I persists at high-$z$. Although the statistics are quite sparse, a substantial increase in either the density and/or the luminosity of this AGN class (at least an order of magnitude increase in active fraction at \\lxbb$\\geq$5$\\times$10$^{42}$ \\ergs\\ when compared to low-$z$) is very similar to the rapid increase recently found in the cluster X-ray AGN population by \\citet{2009ApJ...701...66M}.  However, our study is more complementary to this previous study than confirming because different AGN classes are involved. Optical spectroscopy, primarily derived from the SEXSI study, and deep ground-based and HST imaging for our CRS XPSs finds that these passive AGN are hosted in rather normal, giant elliptical galaxies positioned rather close to the cluster core, whereas the \\citet{2009ApJ...701...66M} AGN are largely emission-line galaxies similar to Seyferts found further out in the clusters near R$_{200}$. Since the energy and mass outflow from Seyferts (i.e., radio-quiet AGN population) is still not well-measured \\citep{2002ApJ...566..699A}, it is not known whether these AGN contribute significantly to ICM heating or not.  Paper I identified the CRS XPSs as low-luminosity BL Lac Objects, suggesting that they possess relativistic jets.  But the absence of radio emission from these sources suggest that these jets are too weak to inject much heat into the ICM (see Paper I for details). However, these CRS X-ray AGN are worthy of further study since deep X-ray imaging \\citep[e.g.,][]{2001AJ....121..662B} finds that a substantial fraction ($\\sim$20\\%) of the faint XPS population is due to passive AGN like these cluster sources. So their detailed nature and cosmological evolution is of considerable interest.  Using the recent scaling law between AGN jet power and the radio luminosity suggested by \\citet{2010ApJ...720.1066C}, P$_{jet}$ $\\propto$ L$_r^{0.7}$, we determine that the excess power from luminous radio galaxies at 23$<$\\logp$<$27 \\radiounits\\ is $\\sim$26$\\times$ higher in \\combozb\\ clusters than in \\lowz\\ clusters, which corresponds to 9--10 fold increase in the AGN jet power if FR~II jets have similar kinetic energies to FR~I jets. Using the more shallow scaling between AGN jet power and the radio luminosity from \\citet{2008ApJ...686..859B} corresponds to a 5--6 fold increase.  Therefore, a larger amount of heat can be deposited into the ICM at high-$z$.  The results presented herein and by previous studies of AGN in high-$z$ clusters \\citep[e.g.,][]{1991ApJ...371...49E,1991ApJ...367....1H,2002AJ....124.1239H,2009ApJ...701...66M} invites new theoretical work on the co-evolution of AGN and the cluster ICM and its effect on ICM heating and cooling.  \\appendix"
    },
    "1107/1107.1056_arXiv.txt": {
        "abstract": "{The morphologies of detected jets in X-ray binaries are almost as diverse as their number. This is due to different jet properties and ambient media that these jets encounter. It is important to understand the physics of these objects and to obtain information about possible sites suitable for particle acceleration in order to explain the observations at very high energies. Here I present the results obtained from the first relativistic hydrodynamical simulations of jets in high-mass microquasars. Our results allow us to make estimates for the emission originated in different sites of the whole structure generated by the jets. These works represent a first step in trying to obtain a deeper understanding of the physics and emission processes related with jets in high-mass microquasars. ",
        "introduction": "High-mass X-ray binaries (HMXB) are the parent population of high-mass microquasars (HMMQ), the latter being HMXB with jet activity. Jets of X-ray binaries (microquasars) are produced close to the compact object (black hole or neutron star) via ejection of material accreted from the stellar companion. Jets in HMMQs propagate in a medium filled by the stellar wind of the massive primary star of the system. At the scales of the binary, the stellar wind can significantly affect the dynamics of the jet, trigger shocks suitable for particle acceleration, and under some conditions the jet may be disrupted by a combination of asymmetric recollimation shocks and hydrodynamical instabilities (see Perucho \\& Bosch-Ramon 2008, Perucho et al. 2010). On the other hand, powerful jets can propagate farther, and once at distances from the injection point much bigger than the separation distance ($z\\gg d_{\\rm orb}$, where $z$ is the axial coordinate and distance to the compact object). Depending on the age of the HMXB, the jet will go through the Supernova remnant (SNR) created by its progenitor or, at later stages, the shocked stellar wind, and then the ISM  (Bosch-Ramon et al. 2011). After crossing these regions, the jet can propagate as far and fast as the inertia of the swept ISM permits, and eventually inflates a wide cocoon that pushes a shell of shocked ISM (Bordas et al. 2009). The occurrence of collisionless shocks in microquasar jets can lead to efficient particle acceleration (e.g., Rieger et al. 2007, Perucho \\& Bosch-Ramon 2008, Bordas et al. 2009, Bosch-Ramon et al. 2011) and non-thermal emission of synchrotron and inverse Compton origin and, possibly, from proton-proton collisions (see Bosch-Ramon \\& Khangulyan 2009 and references therein).  To understand the dynamics of these jets, numerical simulations of their propagation were performed (Peter \\& Eichler 1995; Vel\\'azquez \\& Raga 2000; Perucho \\& Bosch-Ramon 2008, Bordas et al. 2009, Perucho et al. 2010, Bosch-Ramon et al. 2011) together with several analytical treatments (e.g., Heinz \\& Sunyaev 2002; Kaiser et al. 2004; Araudo et al. 2009). In this contribution, I review the results from our simulations of microquasar jets at different scales. ",
        "conclusions": ""
    },
    "1107/1107.2654_arXiv.txt": {
        "abstract": "The prime evidence underpinning the standard $\\Lambda$CDM cosmological model is the CMB power spectrum as observed by \\textit{WMAP} and other microwave experiments. But \\citet{utaneradio} have recently shown that the \\textit{WMAP} CMB power spectrum is highly sensitive to the beam profile of the \\textit{WMAP} telescope. Here, we use the source catalogue from the \\textit{Planck} Early Data Release to test further the \\textit{WMAP} beam profiles. We confirm that stacked beam profiles at Q, V and particularly at W appear wider than expected when compared to the Jupiter beam, normalised either directly to the radio source profiles or using \\textit{Planck} fluxes. The same result is also found based on \\textit{WMAP}-CMBfree source catalogues and NVSS sources. The accuracy of our beam profile measurements is supported by analysis of CMB sky simulations. However the beam profiles from \\textit{WMAP}7 at the W band are narrower than previously found in \\textit{WMAP}5 data and the rejection of the \\textit{WMAP} beam is now only at the $\\approx3\\sigma$ level. We also find that the \\textit{WMAP} source fluxes demonstrate possible non-linearity with \\textit{Planck} fluxes. But including ground-based and \\textit{Planck} data for the bright \\citet{weiland2011} sources may suggest that the discrepancy is a linear offset rather than a non-linearity.  Additionally, we find that the stacked Sunyaev-Zel'dovich (SZ) decrements of $\\approx151$ galaxy clusters observed by \\textit{Planck} are in agreement with the \\textit{WMAP} data. We find that there is no evidence for a \\textit{WMAP} SZ deficit as has previously been reported. In the particular case of Coma we find evidence for the presence of an $\\mathcal{O}(0.1mK)$ downwards CMB fluctuation. We conclude that beam profile systematics can have significant effects on both the amplitude and position of the acoustic peaks, with potentially important implications for cosmology parameter fitting. ",
        "introduction": "\\label{sec:intro}  Cosmic Microwave Background (CMB) experiments such as the Wilkinson Microwave Anisotropy Probe (\\textit{WMAP}) have made significant progress in the study of the primordial temperature fluctuations. Their best fitting power spectra strongly support a spatially flat, $\\Lambda CDM$, universe. This model requires relatively few  parameters, yet apparently manages a compelling concordance between a variety of other cosmological data; SNIa, Large Scale Structure and Big Bang Nucleosynthesis. Although the statistical errors on these power spectra are small, this precision does not necessarily imply accuracy and there remains the potential for systematic errors to alter these conclusions.  Indeed, several anomalies between $\\Lambda CDM$ and the \\textit{WMAP} data have been discussed. Typically these have involved the large-scale temperature multipoles eg: \\citep{bennett2011,liuli}. However, other anomalies in the CMB at smaller scales have also been detected, connected in particular with radio sources \\citep{utaneradio, tommoriond} and SZ decrements from galaxy clusters \\citep{myers2004,bielby2007}  Radio sources are sometimes  regarded as a contaminant in  CMB temperature maps. However, radio point sources prove particularly interesting because they provide a complementary check of the beam measured by the \\textit{WMAP} team from observations of Jupiter \\citep{page2003,hill2009}. Jupiter has a flux of $\\approx$ 1200Jy which is $\\approx3$ orders of magnitude higher than radio source fluxes or CMB fluctuations. This high flux has advantages in terms of defining the wings of the beam profile but has the disadvantage that the calibrating source is much brighter than typical CMB fluctuations. Furthermore, Jupiter only checks the beam on the ecliptic whereas radio sources are spread over the sky. \\citet{utaneradio,tommoriond} made a stacked analysis of radio point sources and found evidence for a wider beam than \\textit{WMAP} measured using Jupiter.  A tentative detection of a non-linear relation between \\textit{WMAP} fluxes and ground based radio telescope fluxes was also found. A thorough analysis of possible systematics did not find an explanation and we return to these issues later in this paper. The beam profile of a CMB telescope  like \\textit{WMAP} is critical because it smoothes the temperature anisotropies and therefore needs to be known accurately to produce the final power spectrum from temperature maps \\citep{page2003,hill2009}.  Various authors have noted small-scale anomalies with respect to the SZ decrements measured by \\textit{WMAP}. SZ decrements are created when CMB photons inverse Compton scatter off hot electrons in galaxy clusters. \\citet{myers2004} first stacked \\textit{WMAP} data at the positions of galaxy clusters and suggested that the profiles were more extended than expected. \\citet{lieu2006} and \\citet{bielby2007} then found that the SZ decrements from \\textit{WMAP} were reduced compared to X-ray predictions, possibly due to the \\textit{WMAP} beam being wider than expected. \\citet{bielby2007} also found that the \\textit{WMAP} decrements were significantly lower than the ground-based SZ measurements by \\citet{bonamente2006} in 38 X-ray luminous clusters.  In their ESZ sample, the \\textit{Planck} team find excellent agreement with the self-similar X-ray estimates of the SZ decrement \\citep{planckxraystats2011,planckszstudy2011}. This is corroborated by the ground based South Pole Telescope Collaboration with their blind SZ selected  cluster sample \\citep{southpole2009}. This compounds the question of why \\textit{WMAP} SZ analyses from \\citet{lieu2006} and \\citet{bielby2007} failed to find such an agreement.  In this paper we use the recent \\textit{Planck} Early Data Release and other radio source data to re-investigate both the \\textit{WMAP} radio source beam profile and SZ anomalies. The \\textit{Planck} Early Release Compact Source catalogue (ERCSC) is of particular interest and provides the basic parameters of radio sources and SZ clusters from the \\textit{Planck} CMB maps. Although, the corresponding temperature maps from which these were estimated have not been released, both radio source fluxes and SZ profile parameters are available as measured by \\textit{Planck}. We can therefore use these to compare \\textit{WMAP} and \\textit{Planck} radio source fluxes directly and also to make  \\textit{WMAP} stacks centred now on the new radio source and SZ cluster lists from \\textit{Planck}. From these stacks, the \\textit{WMAP} beam profile can be inferred and the SZ results from \\textit{WMAP} and \\textit{Planck} compared. Given the higher angular resolution, lower noise and different calibration strategy for \\textit{Planck}, this comparison will allow new insight into the robustness of the \\textit{WMAP} CMB analysis. ",
        "conclusions": "\\label{sec:conclusions}  We have investigated the beam profile of \\textit{WMAP} by comparing beam profiles from radio sources with the Jupiter beam profile. We have compared sources from \\textit{Planck}, \\textit{WMAP}, \\textit{WMAP} CMB-free and NVSS catalogues. We find that in all cases the radio sources show wider profiles than the Jupiter beam with little indication of Eddington bias or dependence on the method of normalisation. Applying our cross-correlation to realistic simulations strongly supports the accuracy of our beam profile measurements. However, it must be said that in the \\textit{WMAP}7 data the W radio source profiles are less wide than previously found  by \\citet{utaneradio} in the \\textit{WMAP}5 release. The rejection of the Jupiter beam is now only $\\approx3\\sigma$ in the \\textit{Planck} radio source comparison. But the rejection of the Jupiter beam is reasonably consistent between the admittedly overlapping radio source samples from \\textit{Planck}, \\textit{WMAP} and NVSS. We have therefore considered explanations for the wide profiles assuming that they are not statistical fluctuations. Two such possibilities are a non-linearity in the \\textit{WMAP} temperature scale and a timing offset in the \\textit{WMAP} scan pattern as discussed by \\citet{liuli}. The narrower profiles measured here compared to the \\citet{utaneradio} \\textit{WMAP}5 profiles increase the possibility of their being explained by a timing offset.  We have also found discrepancies between \\textit{WMAP} fluxes compared to \\textit{Planck} and ground-based fluxes. For $S<30$Jy the \\textit{WMAP} fluxes look to have a non-linear relation with \\textit{Planck} fluxes. However, when the further very bright sources discussed by \\citet{weiland2011} with ground-based and \\textit{Planck} measurements are included then this flux-flux discrepancy appears more like a linear than a non-linear offset.  We have compared stacked \\textit{WMAP} SZ decrements with those measured by \\textit{Planck} and by ground-based observations. In contrast to previous reports we now find \\textit{WMAP} agrees with both the \\textit{Planck} and ground-based data. However this work is not at high enough S/N to distinquish between the timing offset beam of \\citet{SawangwitThesis} and the \\textit{WMAP} Jupiter beam.  We have shown that transforming the Jupiter beam using a model that fits the radio source profiles results in small but significant changes to the \\textit{WMAP} $C_{\\ell}$. At the least, a wider beam would imply a much larger uncertainty in the normalisation and hence the estimate of $\\sigma_8$ from \\textit{WMAP}. Unfortunately, faint radio sources cannot check the \\textit{WMAP} beam at scales larger than $30'$ and a wider beam at these scales could, in principle, change the position, as well as the normalisation, of even the first acoustic peak. Clearly it is important to continue to test the calibration and beam profile of \\textit{WMAP}, particularly in the W band."
    },
    "1107/1107.5597_arXiv.txt": {
        "abstract": "{Radio halos and relics are diffuse radio sources found in galaxy clusters showing significant substructure at X-ray wavelengths . These sources provide important information about  non-thermal processes taking place in the intracluster medium (ICM). Until now only a few dozen relics and halos are known, while models predict that a much larger number of these sources exist. In this paper we present the results of an extensive observing campaign to search for new diffuse radio sources in galaxy clusters. } {The aim of the observations is to create a large sample of diffuse radio sources in galaxy clusters that help to understand the formation of radio relics and halos and can be used to probe the physical conditions of the ICM. } {We carried out radio continuum observations with the Westerbork Synthese Radio Telescope (WSRT),  Giant Metrewave Radio Telescope (GMRT) and Very Large Array (VLA) of clusters with diffuse radio emission visible in NVSS and WENSS survey images. Optical images were taken with the William Herschel and Isaac Newton Telescope (WHT, INT). } {We discovered 6 new radio relics, including a probable double relic system, and 2 radio halos. In addition, we confirm the presence of diffuse radio emission in four galaxy clusters. By constructing a sample of 35 radio relics we find that relics are mostly found along the major axis of the X-ray emission from the ICM, while their orientation is perpendicular to this axis. We also compared the X-ray luminosity and redshift distributions of clusters with relics to an X-ray selected sample from the NORAS and REFLEX surveys. We find tentative evidence for an increase of the cluster's relic fraction with X-ray luminosity and redshift. The major and minor axis ratio distribution of the ICM for clusters with relics is broader than that of the NORAS-REFLEX sample. } {The location and orientation of radio relics with respect to the ICM elongation is consistent with the scenario that relics trace merger shock waves. } ",
        "introduction": "Radio halos and relics are found in massive merging galaxy clusters. These radio sources indicate the presence of magnetic fields and in-situ particle acceleration within the ICM \\citep[e.g.,][]{1977ApJ...212....1J, 2004IJMPD..13.1549G}. Galaxy clusters form through mergers with other clusters and galaxy groups as well as through continuous accretion of gas from the intergalactic medium. Since giant radio halos and relics are found in merging clusters \\citep[e.g.,][]{2001ApJ...553L..15B, 2004ApJ...605..695G, 2007A&A...467...37B, 2009A&A...507..661B, 2010ApJ...721L..82C}, it has been proposed that a small fraction of the energy released during a cluster merger event is channeled into the (re)acceleration of particles.   One scenario for the origin of radio relics is that they trace merger shock waves within the ICM in which particles are accelerated by the diffusive shock acceleration (DSA)  mechanism  \\citep{1977DoSSR.234R1306K, 1977ICRC...11..132A, 1978MNRAS.182..147B, 1978MNRAS.182..443B, 1978ApJ...221L..29B, 1983RPPh...46..973D, 1987PhR...154....1B, 1991SSRv...58..259J, 2001RPPh...64..429M}. In the presence of a magnetic field these particles emit synchrotron radiation at radio wavelengths \\citep[e.g.,][]{1998A&A...332..395E, 2000ApJ...542..608M}. The efficiency with which collisionless shocks can accelerate particles is unknown and may not be enough to produce the observed radio brightness of relics. A closely linked scenario is that of  shock re-acceleration of pre-accelerated electrons in the ICM, which is a more efficient mechanism for weak shocks \\citep[e.g.,][]{2005ApJ...627..733M, 2008A&A...486..347G, 2011ApJ...734...18K}.   An alternative scenario has been proposed by \\cite{2010arXiv1011.0729K} based on a secondary cosmic ray electron model, where the time evolution of magnetic fields and the cosmic ray distribution are taken into account to explain both halos and giant relics. Detailed spectral maps, at multiple frequencies, and measurements of the magnetic field via the polarization properties can test this model. For relics, strong magnetic fields (${B > B_{\\rm{cmb}} = 3.25 \\times (1+z)^2}$) are predicted and the spectral index\\footnote{$F_{\\nu} \\propto \\nu^{\\alpha}$, where $\\alpha$ is the spectral index} at the outer edges of relics should be flat, with $\\alpha \\simeq -1$.   Unlike relics, radio halos are found in the center of merging galaxy clusters and follow the X-ray emission from the ICM. Radio halos have been explained by turbulence injected by recent merger events. The injected turbulence is thought to be capable of re-accelerating relativistic particles \\citep[e.g., ][]{2001MNRAS.320..365B, 2001ApJ...557..560P, 2005MNRAS.357.1313C}. Alternatively, the energetic electrons are secondary products which originate from hadronic collisions between relativistic protons and thermal ions \\citep[e.g.,][]{1980ApJ...239L..93D, 1999APh....12..169B, 2000A&A...362..151D, 2010ApJ...722..737K, 2011A&A...527A..99E}.  Recent observations put tension on the secondary models \\citep[e.g, ][]{2010MNRAS.401...47D, 2010MNRAS.407.1565D, 2011ApJ...728...53J, 2011MNRAS.412....2B, 2011A&A...530A..24B}. The existence of ultra-steep spectrum radio halos is also claimed not to be in agreement with these secondary models \\citep{2008Natur.455..944B}, but they can be explained by the turbulent re-acceleration model. However, currently only a few of these ultra-steep spectrum radio halos are known so more observations, such as presented in this paper, are needed to increase this number and provide better measurements of the radio spectra.  In the last decade a number of successful searches have been carried out to find new diffuse radio sources in galaxy clusters \\citep[e.g.,][]{1999NewA....4..141G, 2000NewA....5..335G, 2001A&A...376..803G, 2001ApJ...548..639K, 2003A&A...400..465B, 2007A&A...463..937V, 2008A&A...484..327V,2009A&A...507.1257G, 2009A&A...508...75V, 2009ApJ...697.1341R}. However, our understanding of the formation of these sources is still limited. Models for the formation of relics and halos can be tested through statistical studies of correlations between the non-thermal radio emission and global properties of the clusters, such as mass, temperature, and dynamical state \\citep{2000ApJ...544..686L,2006MNRAS.368..544F,2006MNRAS.369.1577C, 2007MNRAS.378.1565C, 2008A&A...480..687C, 2010A&A...509A..68C}.  We recently discovered two large radio relics in the NVSS and WENSS surveys \\citep{2010Sci...330..347V,2011A&A...528A..38V}. Interestingly, these relics remained unnoticed for about 15 years. This suggests that more diffuse radio sources could be discovered by inspection of the NVSS \\citep{1998AJ....115.1693C} and WENSS \\citep{1997A&AS..124..259R} survey images. We therefore carried out a visual inspection of the NVSS and WENSS images around known clusters detected by ROSAT \\citep{1999A&A...349..389V, 2000IAUC.7432R...1V, 1998MNRAS.301..881E, 2000ApJS..129..435B, 2007ApJ...662..224K, 2002ApJ...580..774E, 2004A&A...425..367B}.  Candidate clusters hosting diffuse radio emission were observed with the WSRT, GMRT and/or VLA. Clusters with existing published observations were not re-observed. In this paper we present the radio images and global properties of the clusters. In addition, we investigate the position and orientation of radio relics with respect the ICM, and compare the X-ray luminosity and redshift distributions of clusters with relics to an X-ray selected sample. In a follow-up paper we will present the polarization and detailed spectral properties of the radio emission in these clusters. The layout of this paper is as follows. In Sect.~\\ref{sec:obs-reduction} we give an overview of the observations and the data reduction. In Sect.~\\ref{sec:results} we present the radio images and give an short overview of the cluster's properties. We end with discussions and conclusions in Sects.~\\ref{sec:discussion} and \\ref{sec:conclusion}.   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$. All images are in the J2000 coordinate system.  \\begin{table*} \\begin{center} \\caption{Observations} \\begin{tabular}{lllllll} \\hline \\hline Cluster &observation date & frequency & bandwidth   & integration time & map resolution & rms noise\\\\ &                             & MHz          & MHz           &  hr                         &   arcsec              &  $\\mu$Jy~beam$^{-1}$\\\\ \\hline Abell 1612  &  GMRT, 13 May, 2009 & 325  & 32 & 4.0& $11.6\\arcsec~\\times9.5\\arcsec$ & 236\\\\ & GMRT, 22 Nov, 2009 & 610, 241 & 32, 6 & 2.5, 2.5 &  $7.7\\arcsec~\\times4.7\\arcsec$, $21\\arcsec \\times 12\\arcsec$& 64, 777\\\\ Abell 523    & VLA Aug, 25, 2005;&  1425 & 37.5 & 1.1$^{a1}$, 3.7$^{a2}$ & $21\\arcsec~\\times20\\arcsec^{d}$& 40\\\\ & Nov 28, 2005; Dec 28, 2009                                                                       & \\\\ % Abell 697$^{e}$  &   WSRT, 24 Aug, 2009 & 1382, 1714& 160, 160& 6.0, 6.0&$34\\arcsec~\\times21\\arcsec^{b}$ & 24, 32\\\\ Abell 3365$^{e}$  & WSRT, 22 Feb, 2009 & 1382& 160& 12.0&$108\\arcsec~\\times13\\arcsec$&  29\\\\  Abell 3365  & VLA, 30 Sep, 2009   & 1425& $2 \\times 50$ & 0.7$^{g1}$, 2.5$^{g2}$ & $43\\arcsec~\\times~35\\arcsec$, $13.5\\arcsec~\\times~ 9.2\\arcsec$ & 84, 27 \\\\ &  30 May, 2009    &  & \\\\ Abell 746$^{e}$    & WSRT, 19 Sep, 2009 & 1382& 160 & 6.0&$23\\arcsec~\\times18\\arcsec$ & 28\\\\ Abell 2034$^{f}$ & WSRT, 26 Jul, 2009 & 1382& 160 & 6.0&$30\\arcsec~\\times16\\arcsec$ & 24\\\\ Abell 2061$^{f}$ & WSRT, 23 Jul, 2009& 1382, 1714& 160, 160 & 6.0, 6.0&$32\\arcsec~\\times16\\arcsec^{b}$ & 22, 26\\\\ RXC J1053.7+5452$^{e}$ & WSRT, 14 Mar, 2009 & 1382& 160 & 6.0 &$24\\arcsec~\\times18\\arcsec$ & 30\\\\ % CIZA J0649.3+1801  & GMRT, 22 Nov, 2009 & 610, 241& 32, 6  & 3.0, 3.0& $25\\arcsec~\\times25\\arcsec^{c}$, $17\\arcsec \\times 14\\arcsec$& 515, 1800\\\\ RX J0107.8+5408$^{e}$ & WSRT, 29 Aug, 2009& 1382 & 160 & 6.0& $21\\arcsec~\\times17\\arcsec$ & 29\\\\ \\hline \\end{tabular} \\label{tab:observations} \\end{center} $^{a1}$ D array\\\\ $^{a2}$ C array\\\\ $^{b}$ the WSRT 1714~MHz image was convolved to the same resolution as the 1382~MHz image\\\\ $^{c}$ convolved to a beam of  $25\\arcsec~\\times25\\arcsec$\\\\ $^{d}$ D and C array data were combined\\\\ $^{e}$ The minimum baseline length is 36~m. \\\\ $^{f}$ The minimum baseline length is 54~m. \\\\ $^{g1}$ DnC array\\\\ $^{g2}$ CnB array\\\\ \\end{table*} ",
        "conclusions": "We have presented WSRT, VLA and GMRT observations of galaxy clusters with diffuse emission selected from the WENSS and NVSS surveys. We find  peripheral radio relics in the clusters Abell 1612, Abell 746 and CIZA~J0649.3+1801 and a smaller relic Abell~2034. Abell~3365 seems to host a double radio relic system. Our observations reveal radio halos in the clusters Abell~746 and RX~J0107.8+5408. We confirm the presence of radio relics in Abell~2061, RXC~J1053.7+5452 and diffuse emission in Abell~523 and Abell~2034 (for which the classification is uncertain). In addition, we detect the radio halo Abell~697 providing additional flux measurements around 1.4~GHz.  By constructing a sample of 35 relics we have found that relics are generally located along the major axis (which can be used as a proxy for the merger axis) of the cluster's elongated ICM. Their orientation is mostly perpendicular to this major axis. The distribution of the major-minor axis ratios for the relic cluster sample is broader than that of the NORAS-REFLEX sample. These results are consistent with the scenario that relics trace shock waves which form along the merger axis of clusters. We also compared the redshift and X-ray luminosity distributions of clusters with relics to a sample of clusters from the NORAS and REFLEX surveys. Interestingly, we find indications that the observed fraction of clusters hosting relics increases with X-ray luminosity and redshift. However, selection biases for radio relics play a role and this needs to be investigated further. The significantly improved sensitivity of upcoming radio telescopes (e.g., the EVLA and LOFAR) will allow to find many more radio relics. Correlating this larger relic sample with cluster samples selected at X-rays wavelengths will provide a powerful tool to constrain the evolution of relics in galaxy clusters.    \\label{sec:conclusion}"
    },
    "1107/1107.5845_arXiv.txt": {
        "abstract": "The exact composition of a specific class of compact stars, historically referred to as ``neutron stars'', is still quite unknown. Possibilities ranging from hadronic to quark degrees of freedom, including self-bound versions of the latter have been proposed. We specifically address the suitability of strange star models (including pairing interactions) in this work, in the light of new measurements available for four compact stars. The analysis shows that these data might be explained by such an exotic equation of state, actually selecting a small window in parameter space, but still new precise measurements and also further theoretical developments are needed to settle the subject. ",
        "introduction": "The first reasonable ideas about the composition of compact stars stated that matter under extreme densities is composed of hadrons, i.e., mainly neutrons with small fractions of protons and also electrons. Further theoretical developments and modern experimental results opened the window to other possibilities as the existence of hyperons inside such stars due to the high densities, and a whole new class of condensates (see, for example, \\cite{Weber} and references therein). Another possibility arisen in the '70s is the existence of deconfined quarks inside compact objects, either located only in the inner core of these stars or present up to their surfaces. The latter extreme possibility was suggested about three decades ago \\cite{Itoh, Bodmer, Chin, Terazawa, Wit}, emerging as an astrophysical realization of the stability scenario of the so-called strange matter.  Consequences of the so-called {\\it strange matter hypothesis} for compact stars have been extensively analyzed within the simple framework of the MIT Bag Model \\cite{Farhi84, AFO} and the Nambu-Jona-Lasinio model \\cite{Alford, BenvLug98, Yin08} (see \\cite{Buballa05} for an overview on the subject). After an initial round of perturbative considerations \\cite{BailinLove, Horvathetal91}, the possibility of non-perturbative pairing between quarks \\cite{Alford, Rapp, Alford01, RajWilc, Rajagopal2001, German} allowed a new view on this subject due to a great enhancement on the window of stability for strange quark matter (having the strange quark mass, bag constant and the novel pairing energy of the quark condensate as parameters in this approach). This state of totally paired 3-flavor quarks is called the color-flavor locked (CFL) strange quark matter, in which quarks form Cooper pairs, thus lowering further the energy of the system.  Much work has been put forward in order to have a better understanding on the actual composition of neutron stars, both observational and theoretical. Nevertheless, there is no absolute conclusion for any model as yet.  We intend in this work to contribute to the understanding of how the new and much more precise astrophysical measurements of the mass and radius of neutron stars \\cite{4U1608, EXO1745, 4U1820, PSRJ1614} can help revealing the viability of exotic quark star models. We analyze in particular the viability of the class of CFL strange quark matter models. They are likely to be the most favorable candidates for self-bound stars, in spite that more complex models can be devised (for example intermediate phases such as LOFF may be present \\cite{LOFF1, LOFF2, Casalbuoni04}), and they would constitute a benchmark in testing exotic compact star composition.  In the rest of this work we quantitatively show that accurate astrophysical measurements can already constrain important parameters of QCD, like the Cooper pair energy gap ($\\Delta$) and the strange quark mass ($m_{s}$), and show how this is done given a subset of the present data. ",
        "conclusions": "Precise measurements for masses and radii of compact stars are very important for constraining their composition. However, besides being precise, it is necessary that they also be accurate for a clear comparison of the mass-radius relation coming from very different compositions with data.  Given the present controversy of the radius determination quoted in the paper, this appears to be the main issue to be resolved for an advance of the conclusions of studies like ours. Meantime we can only conclude that only with small error bars for both mass and radii determinations (as in the analysis of \\cite{4U1608, EXO1745, 4U1820}) the parameters of the CFL equation of state can be effectively constrained. Otherwise, several sets of parameters defining the CFL EoS can be used to match the data. In this case, more advanced detection techniques and physical analysis are required to constrain the radius $R_{ph}$ in X-ray burst models.  Since one should {\\it not} treat $\\alpha$ and $B_{eff}$ as independent quantities in a parameterized EoS for CFL strange quark matter like the one proposed here, an analysis performed with a parameterization with correlated coefficients should be explored. We intend to address this subject in a future work.  An additional interesting report is the very low mass of at least one neutron star mass observed ($M=0.87\\pm 0.07 M_{\\odot}$ for 4U 1538-52) \\cite{Rawls}, actually not well explained by ordinary stellar evolution arguments. This ``low-end'' mass extreme is also important for the quest of exotic compositions in their interiors. In fact, confirming such a measurement would imply a radius of $R \\leq 9$ km within the self-bound hypothesis, while the measurement of, say, $R \\geq 11$ km would be totally inconsistent with it, pointing to a normal composition. We are at the edge of finally addressing the fundamental composition of compact stars, at least on a first level that leaves a handful of possibilities left out of a huge number of proposals."
    },
    "1107/1107.3160.txt": {
        "abstract": "With the availability of the huge amounts of data produced by current and future large multi-band photometric surveys, photometric redshifts have become a crucial tool for extragalactic astronomy and cosmology. In this paper we present a novel method, called Weak Gated Experts (WGE), which allows to derive photometric redshifts through a combination of data mining techniques. %With respect %to the many other existing methods published in literature, the WGE offers the advantage that it requires little %or no fine-tuning and it can be applied to different classes of extragalactic sources (galaxies and quasars). %and is scalable to very large data sets. \\noindent The WGE, like many other machine learning techniques, is based on the exploitation of a spectroscopic knowledge base composed by sources for which a spectroscopic value of the redshift is available. This method achieves a variance $\\sigma^2(\\Delta z)\\!=\\!2.3\\!\\cdot\\!10^{-4}$ ($\\sigma^2(\\Delta z)\\! =\\!0.08$, where $\\Delta z = x_{\\mathrm{phot}} - z_{\\mathrm{spec}}$) for the reconstruction of the photometric redshifts for the optical galaxies from the SDSS and for the optical quasars respectively, while the Root Mean Square (RMS) of the $\\Delta z$ variable distributions for the two experiments is respectively equal to 0.021 and 0.35. The WGE provides also a mechanism for the estimation of the accuracy of each photometric redshift. We also present and discuss the catalogs obtained for the optical SDSS galaxies, for the optical candidate quasars extracted from the DR7 SDSS photometric dataset\\footnote{The sample of SDSS sources on which the accuracy of the reconstruction has been assessed is composed of bright sources, for a subset of which spectroscopic redshifts have been measured.}, and for optical SDSS candidate quasars observed by GALEX in the UV range. The WGE method exploits the new technological paradigm provided by the Virtual Observatory and the emerging field of Astroinformatics. ",
        "introduction": "\\label{sec:introduction}  The ever growing amount of astronomical data provided by the new large scale digital surveys in a wide range of the EM spectrum has been challenging the way astronomers carry out their everyday analysis of astronomical sources. These new data sets, for their sheer size and complexity, have extended beyond the human ability to visualize and correlate complex data, thus triggering the birth of the new technological approach and methodology which is often labeled as ``astroinformatics\", a new discipline which lies at the intersection of many others: data mining, parallel and distributed computing, advanced visualization, web 2.0 technology, etc. \\cite{borne2009,ball2010}. X-informatics (where the X stands for any data rich discipline), is growingly being recognized as the fourth leg of scientific research after experiment, theory and simulations (see \\emph{The Fourth Paradigm}, \\cite{hey2009}). In this paper we shall present a new method for the estimation of photometric redshifts which fully take advantage of many of these new methodologies.  \\noindent In the past, for many tasks such as, for instance, classifying different types of sources, determining the redshifts of galaxies, etc. astronomers had to rely mainly on spectroscopic observations which are still very demanding in terms of precious telescope time.  Even though spectroscopy is still fundamental to gain insights into many physical processes, the unprecedented abundance of accurate photometric observations for very large samples of sources, has led to the development of what we can call {\\it candidates astronomy}, i.e. the branch of astronomy which exploits photometry to accomplish tasks which in the past would have required spectroscopic data. This discipline stems from a long and rich tradition of astronomical techniques based on the use of photometric information in low dimensional \\emph{features} space (for example, colour-colour selection techniques). The main differences relative to these classical methodologies reside in the statistical techniques and the size of the dataset considered (in terms of both the number of members and the dimensionality of the datasets). In those cases where a very accurate evaluation of the uncertainties affecting the estimate is possible, the loss of accuracy and effectiveness which is implicit in candidates astronomy, is compensated by the possibility to obtain very extensive samples with limited effort. In the last few years, candidates astronomy has found many applications, such as, for instance, the determination of the spatial distribution of visible matter on very large scales through photometric redshifts \\cite{arnaltemur2009}. In such cases, the statistical tools used to characterize the description of the distribution of the sources are specifically designed to trade-off between the lower accuracy of the derived quantities (e.g. photometric redshifts with respect to spectroscopic ones) and the increased statistics arising from the significantly larger size of the samples of sources. Another example is the study of the distribution of quasars through the use of photometric redshifts and reliable catalogs of candidate quasars selected on the basis of their photometric properties rather than through spectroscopic confirmations. The advantages of candidates astronomy over traditional astronomy are obvious: for instance, in the latter case, the number of quasars selected via spectroscopic identification in the Sloan Digital Sky Survey (SDSS) Data Release 7 (DR7) is $\\sim7.5\\!\\cdot\\!10^4$ \\cite{abazajian2009}, while the number of candidate quasars extracted from photometric data with effective methods involving the modeling of the distribution of sources in the color space is almost an order of magnitude larger, ranging from $\\sim\\!10^6$ found by \\cite{richards2009} to the higher $\\sim\\!2.1\\cdot\\!10^6$ in \\cite{dabrusco2009}, the latter figures being much closer to the theoretically predicted number of quasars expected to lie within the limiting flux of the SDSS, $\\sim\\!1.3\\!\\cdot\\!10^6$ reported in \\cite{richards2009}.  Photometric redshifts are important for a large spectrum of cosmological applications, such as, to quote just a few: weak lensing studies of galaxy clusters \\cite{abdalla2008}, the determination of the galaxy luminosity function \\cite{subbarao1996}, studies of specific types of cosmic structures like, for instance, the photometric redshifts derived in \\cite{dabrusco2007} which were used to investigate the physical reality of the so-called Shakhbazhian groups, to derive their physical characteristics as well as their relations with other galaxy structures of different compactness and richness \\cite{capozzi2009}.  Many different methods for the evaluation of  photometric redshifts are available in literature. Without entering into much detail, it is worth reminding that all methods are based on the interpolation of \\emph{a priori} knowledge available for more or less large sets of templates and differ among themselves only in one or both of the following aspects: i) the way in which the \\emph{a priori} Knowledge Base (KB, for a detailed definition of KB, see section \\ref{sec:datamining}) is constructed (higher accuracy spectroscopic redshifts or empirically or theoretically derived Spectral Energy Distributions (hereafter SEDs), and ii) the interpolation algorithm or method employed. In this context, modern wide-field mixed surveys combining multi-band photometry and fiber-based spectroscopy and thus providing both photometric data for a very large number of objects and spectroscopic information for a smaller but still significant subsample of the same population, provide all the information needed to constrain the fit of an interpolating function mapping the space of the photometric features. Most if not all photometric redshifts methods have been tested on the Sloan Digital Sky Survey (SDSS) which is a remarkable example of these ``mixed surveys\", which has allowed noticeable advancements in the field of extragalactic astronomy and, over the years, has also become a sort of standard benchmark to evaluate performances and biases of different methods. Nonetheless, it should be noticed that the SDSS spectroscopic sample is not unbiased, since limited to a bright subset of galaxies and quasars observed in the optical range and selected according to spectroscopic methods. The peculiar characteristics of the source samples for which both photometric and spectroscopic measurements are available should always borne in mind when considering the effectiveness of the machine learning methods tested.  \\noindent As it will become evident in section \\ref{sec:wge}, one of the main problems encountered in evaluating photometric redshifts is the critical dependence of the final accuracy on the parameters needed to fine-tune the method and the nature of the sources (i.e., galaxies or quasars). For example, in the template fitting methods, part of the degeneracy between the spectroscopic redshift and colors of the sources can be minimized by a wise choice of the SED templates \\cite{bruzual2010}, at the cost of introducing biases in the final estimates of the photometric redshifts. In other data mining applications the same degeneracy can be minimized by applying priors derived from the distribution of the spectroscopic redshifts for the sources belonging to the KB, like in \\cite{dabrusco2007}.  In what follows, we shall just summarize some aspects which appear to be relevant for the class of the interpolative methods. Such methods differ in the way the interpolation is performed, and the main source of uncertainty is the fact that the fitting function is just an approximation of a more complex and unknown relation (if any) existing between colors and the redshift (for example, see \\cite{csabai2003}). Moreover, due to different observational effects and emission mechanisms, a single approximation can hold only in a given range of redshifts or in a limited region of the \\emph{features} space \\cite{dabrusco2007}. In the last few years, in order to overcome the effects of the oversimplification of the relation between observables and spectroscopic redshifts, several methods based on statistical techniques for pattern recognition aimed at the accurate reconstruction of the photometric redshifts for both galaxies and quasars have been developed (and in most cases applied to SDSS data): polynomial fitting \\cite{connolly1995, budavari2001, li2008}, nearest neighbors \\cite{csabai2003, ball2008, budavari2001}, neural networks \\cite{firth2003,collister2004,vanzella2004,collister2007,dabrusco2007,yeche2010}, support vector machines \\cite{wadadekar2005}, regression trees \\cite{carliles2010}, Gaussian processes \\cite{way2006,bonfield2010} and diffusions maps \\cite{freeman2009}.  \\noindent These methods, when applied to SDSS galaxies in the local universe (i.e. $z\\!<\\!0.5$), lead to similar results, with a dispersion $\\mathbf{RMS(\\Delta z)\\!\\sim\\!0.02}$. The extension of these methods to the intermediate redshifts range ($z\\!<\\!0.8$) is in theory possible for both quasars and galaxies and for the brightest sources in the SDSS dataset, by adding near infrared photometry to the Sloan optical photometry and by using as KB the large SDSS spectroscopic sample, sometimes combined with redshifts measured in other deeper surveys like, for instance, the 2SLAQ \\cite{croom2004}. Even though in this redshift range the estimated photometric redshifts seem not to be affected by any peculiar systematic effect, all these methods suffer from strong degeneracies in specific regions of the photometric \\emph{features} space when applied to sources, like quasars, which can be found at larger redshifts and whose spectra typically present strong emission and absorption features, because of several different effects, often depending on the specific method used: reduced statistics, strong evolutionary effects or observational effects, like peculiar spectroscopic features being shifted off and in the photometric filters adopted in a specific instrument (like shown for SDSS quasars in \\cite{ball2008}).  \\noindent Such degeneracies manifest themselves mainly through high local fractions of catastrophic outliers, i.e. sources with photometric redshifts estimates differing dramatically from the spectroscopic value. Some of the previously mentioned techniques address the problem of the catastrophic outliers by providing probabilistic estimates of the photometric redshifts \\cite{ball2008}, at the cost of an increased computational burden, which may lead to an overall worse scalability. The Weak Gated Experts method (hereafter WGE) described in this paper %, whose description and the implementation used for the determination of the photometric redshifts are %provided in sections \\ref{sec:datamining} and \\ref{sec:wgeimplementation}, respectively, has been designed to be accurate, relatively fast when compared to the other approaches available in literature, and easily scalable in order to allow the processing of very large throughputs (like those that will be produced by the large synoptic surveys of the future). As it will be discussed in what follows, the WGE method is general and comprehensive since it adapts to different types of sources without requiring a specific fine tuning. The WGE is the second step of an automated machine learning method whose ultimate goal is to ease the exploitation of ongoing and planned multi-band extragalactic surveys. While not completely removing the catastrophic outliers (task which is impossible to achieve due to the physical limitations, as mentioned above), WGE achieves a fair characterization of the regions of the photometric feature space where the degeneracies happen and is consistently able, as discussed in section \\ref{sec:errors}, to flag the photometric redshifts values which most likely are catastrophic outliers. It is worth noticing that in our analysis we never take into account possible time variability as it is done, for instance, in \\cite{salvato2009}.  \\noindent The paper is structured as follows: in section \\ref{sec:datamining} the general features and design principles of the WGE method are discussed. In section \\ref{sec:wge}, more details of the specific implementation of the WGE used for the problem of photometric redshift reconstruction and of the algorithms employed are provided. The description of the datasets used for the experiments\\footnote{Throughout this paper, the word experiment will refer to a complete run of the WGE method, as customary in the data mining jargon.} and the \\emph{feature} selection criteria can be found in section \\ref{sec:boks}, while the experiments for the determination of photometric redshifts for galaxies and quasars are described in section \\ref{sec:experiments}. The final catalogs of photometric redshifts for the SDSS galaxies and candidate quasars are presented in sections \\ref{sec:catgal}, \\ref{sec:catqua} and \\ref{sec:catquauv} respectively, together with details on the distribution of the catalogs to the community. A thorough discussion of the performances of the method for the reconstruction of the photometric redshifts can be found in section \\ref{sec:accuracy}, while the determination of the errors on the photometric redshifts estimates and the discussion of the catastrophic outliers are described in section \\ref{sec:errors}. The conclusion and a summary of the results can be found in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  %In this paper we have presented the WGE, an original method for the determination of the photometric redshifts for all kind of %extragalactic sources, based on the combination of clustering and regression techniques. This method is able to resolve most %of the degeneracies in the distribution of the KB in the \\emph{features} space and to derive accurate estimates of both the %photometric redshifts values and the errors on the $z_{\\mathrm{phot}}$. The WGE has been then applied to the problems of the determination %of photometric redshifts of optical galaxies and to the quasars extracted from the SDSS-DR7 database. The accuracy of the %reconstruction of the redshifts for optical galaxies is $\\left\\langle \\Delta z \\right\\rangle = 0.015$, with $\\sim86.9\\%$ of sources %with $\\Delta z < 0.03$. On the other hand, the same diagnostics for the estimation of $z_{\\mathrm{phot}}$ for candidate quasars %are $\\left\\langle \\Delta z \\right\\rangle = 0.21$ and $\\%(\\Delta z\\!> \\!0.3 = 80)$. %Another advantage of the WGE over most methods in the literature is its scalability, i.e. the ability to process very large %datasets in a short time, which is increasingly one of the requirements of data mining techniques in order to be ready for %the future optical synoptic surveys (PanStars, LSST), that are poised to produce overnight an amount of data similar o %larger than the total amount of data produced by the whole SDSS survey. A thorough discussion of the implementation %of the method and a comparison with several other methods for the evaluation of the photometric redshifts applied to similar data %are also present in the paper. The results of the comparison a large set of statistical diagnostics %shows that the WGE performs better or similarly than all the other methods, when applied to similar datasets. %The results of the best experiments with the WGE for both optical galaxies and quasars have been used to produce %the catalogs of photometric redshifts of $\\sim 3.2\\!\\cdot\\!10^7$ galaxies photometrically selected and of a sample of $\\sim 2.1\\!\\cdot\\!10^6$ %optical candidate quasars from \\cite{dabrusco2009} respectively. Both catalogs are publicly available (see \\ref{sec:catgal} and %\\ref{sec:catqua} for details). In this paper, we have also shown the results of the application of the WGE method to a relatively %small sample of spectroscopically selected optical SDSS quasars with ultraviolet (GALEX) photometry available, in order to %illustrate how the accuracy of the determination of the photometric redshifts will greatly benefit from the availability of %multi-wavelength photometry. % %QUI VANNO MESSI ANCHE I VANTAGGI DEL WGE RELATIVI ALLA SCALABILITA', %ALLA POSSIBILITA' DI RIADDESTRARE IL METODO  OVERNIGHT ETC. ETC. COME DISCUSSO CON %OMAR - VEDI  VERSIONE STAMPATA CHE HA OMAR.  The Weak Gated Expert or WGE is an original method for the determination of the photometric redshifts capable of working on both galaxies and quasars. The WGE, which is based on a combination of clustering and regression techniques, is able to mitigate most of the degeneracies which arise from the distribution of KB templates in the \\emph{features} space, and to derive accurate estimates of both the photometric redshifts values and of their errors. Besides giving a detailed description of how the WGE works, in this paper we have also presented an application of the WGE to the determination of photometric redshifts of optical galaxies and to the candidate quasars with optical and ultraviolet photometry, both extracted from the SDSS-DR7 database. The accuracy of the reconstruction of the redshifts for optical galaxies, obtained by comparing photometric and spectroscopic redshifts, can be expressed a robust estimate of the dispersion of the $\\Delta z$ variable, which is equal to MAD($\\Delta z$) = 0.011 with $\\sim86.9\\%$ of the sources within $\\Delta z_{3}$. The same diagnostics for the estimation of $z_{\\mathrm{phot}}$ for candidate quasars are MAD($\\Delta z$) = 0.11 and $\\%(\\Delta z_{3}) = 80.5$ when only optical photometry is used, reaching MAD($\\Delta z$) = 0.061 and $\\%(\\Delta z_{3}) = 91.4$ when the photometric redshifts are evaluated using both optical and ultraviolet photometry. %Another advantage of the WGE over most methods in the literature is its scalability, i.e. the ability to process very large %datasets in a short time. A thorough discussion and a comparison of the WGE with several other methods applied to the same or similar data is also provided in the paper. To perform such comparison, a large set of statistical diagnostics shows that the WGE performs better than or similarly to all the other methods. The results of the best experiments with the WGE for optical galaxies and quasars have been used to produce the catalogs of photometric redshifts of $\\sim 3.2\\!\\cdot\\!10^7$ galaxies photometrically selected, a sample of $\\sim 2.1\\!\\cdot\\!10^6$ optical candidate quasars from \\cite{dabrusco2009} with photometric redshifts estimated using optical only photometry and a smaller catalog of more accurate photometric redshifts derived from optical and ultraviolet photometry for a subset of $\\sim 1.6\\!\\cdot\\!10^5$ optical candidate quasars respectively. All catalogs will be publicly available and a complete description of the parameters associated to each photometric redshift estimates is available (see \\ref{sec:catgal}, \\ref{sec:catqua} and \\ref{sec:catquauv} respectively for details on the catalogs). In this paper, we have also shown the results of the application of the WGE method to a relatively small sample of spectroscopically selected optical SDSS quasars for which also the ultraviolet (GALEX) photometry was available. %For what scalability is concerned, we wish to stress that for the WGE, as well as for most other supervised DM methods, %one further advantage is the ability to generalize the method. Since the largest computational load is in the training phase, once the WGE has been trained and has has achieved the required accuracy (either by matching some a priori constraint or by convergence), it can be ``frozen\" and newly acquired data falling in the same region of the \\emph{features} space sampled by the KB can be processed without the need for a re-training of the method. This implies that, regardless the rate at which data are acquired, the WGE can produce estimates of photometric redshifts in real-time. If needed, a new training of the method can be performed off-line when a larger/improved KB becomes available. This requirement is becoming of the utmost importance for data mining techniques in order for them to cope with the data streams foreseen for the current and future optical synoptic surveys (such as Pan-STARRS or the LSST) that will produce overnight an amount of data (images and catalog) similar or even larger than the total amount of data collected by the SDSS. \\noindent It is worth stressing that the WGE is part of the larger realms of Astroinformatics and Data Mining. As a data-driven discipline, through the application of Data Mining methods, Astroinformatics can provide Astronomy and Astrophysics with a framework for tackling new problems or old problems with a novel approach: in particular, where the traditional approach uses data from observations in order to prove or disprove an hypothesis, with Data Mining we want data itself to provide hypotheses that can be then proved or disproved with more accurate follow-up observation. For example, using a catalog of photometric redshifts for galaxies of the SDSS DR7 survey, \\cite{capozzi2009} put constraints on the nature of the so called Shakbazian groups by studying the properties of such groups as they appeared to be in the de-projected space. %Later follow-ups (reference?) have confirmed their hypotheses. This data-driven approach is well described also by the fact that by using machine learning methods many assumptions can be dropped in favor of a more agnostic approach: for instance, by employing machine learning techniques to the photometric redshift problem, one can drop any assumptions on the form of the SED of the source, so that it is up to the model, for example a neural network, to find a representation of the highly non-linear relation between the photometric information and the spectroscopic redshift, instead of fitting the data with a set of SED templates. However, since the hypothesis driven approach of template fitting has noticeable advantages, it can be useful to underline an interesting feature offered by the WGE: it is possible, through the WGE, to link together different \\emph{experts} employed in complex architectures, in which different predictors can be integrated to take advantage of the peculiar strengths of each of them. Even an algorithm which does not belong to the domain of the DM techniques could be consistently used together with machine learning experts. In this case, however, the predictors not based on DM techniques will not be trained in the first step of the training algorithm and will only participate to the training of the gate predictor. This feature was not exploited in this work, but a likely outcome of such a hybrid approach will be the creation of mixed WGE architectures in which empirical machine learning algorithms cooperate with more traditional algorithms based on physical knowledge, for instance neural networks and SED template fitting\\footnote{For a review of the most used template fitting methods in the literature see \\cite{hildebrandt2010}).}  \\noindent One interesting feature of this approach is the generalization offered by the WGE: linking together different \\emph{experts} can lead to complex architectures in which different predictors can be integrated to take advantage of the peculiar strengths of each of them. Even an algorithm which does not belong to the domain of the DM techniques can be consistently used together with machine learning experts. In this case, however, the predictors not based on DM techniques will not be trained in the first step of the training algorithm and will only participate to the training of the gate predictor. A likely outcome of such hybrid approach will be the creation of mixed WGE architectures in which empirical machine learning algorithms cooperate with more classical algorithms based on physical consideration or models typical of the specific domain. For instance, for the particular problem of the estimation of photometric redshifts, the WGE method could be used to integrate machine learning algorithms and the empirical methods based on SED template fitting (for a review of the most used template fitting methods in the literature see \\cite{hildebrandt2010}). %The WGE method will be made available to the community through a specific WEB application %accessible at the DAME (DAta Mining \\& Exploration) project website \\cite{DAME}."
    },
    "1107/1107.5132_arXiv.txt": {
        "abstract": "The mass-dependent equilibrium stable isotope fractionation between different materials is an important geochemical process. Here we present an efficient method to compute the isotope fractionation between complex minerals and fluids at high pressure, $P$, and temperature, $T$, representative for the Earth's crust and mantle. The method is tested by computation of the equilibrium fractionation of lithium isotopes between aqueous fluids and various Li bearing minerals such as staurolite, spodumene and mica. We are able to correctly predict the direction of the isotope fractionation as observed in the experiments. On the quantitative level the computed fractionation factors agree within $1.0\\,\\permil$ with the experimental values indicating predictive power of {\\it ab initio} methods. We show that with {\\it ab initio} methods we are able to investigate the underlying mechanisms driving the equilibrium isotope fractionation process, such as coordination of the fractionating elements, their bond strengths to the neighboring atoms, compression of fluids and thermal expansion of solids. This gives valuable insight into the processes governing the isotope fractionation mechanisms on the atomic scale. The method is applicable to any state and does not require different treatment of crystals and fluids. ",
        "introduction": "\\label{} The fractionation of stable isotopes between various materials is of importance in geoscience, as the variation in isotope content provides valuable information on processes and interaction between atmosphere, biosphere, geosphere and hydrosphere. Although there is substantial analytical work performed in this area, reliable computational methods to predict isotope fractionation factors have been available only recently, proving that they can contribute towards understanding geochemical mechanisms responsible for production of isotope signatures \\citep{D97,YMMW01,S04,DK08,HS08,ML09,SMH09,Z09,HS10,RC10,RB10,Z09,Z10}.  \\begin{table*}[t] \\caption{Lattice parameters of the investigated Li-bearing silicates. $N_{atoms}$ is the number of atoms in the modeled supercell. The units are $\\rm \\AA$ and degrees.} \\label{T1} \\centering \\begin{tabular}{lcccccc}     % \\hline\\hline & staurolite & spodumene & mica 1M & mica 2M1 & mica 2M2 & mica 3T \\\\ \\hline a &  7.848 &  9.463  & 5.20  & 5.209  & 9.04  & 5.200  \\\\ b & 16.580 &  8.392  & 9.01  & 9.053  & 5.22  & 5.200  \\\\ c &  5.641 & 10.436  & 10.09  & 20.053  & 20.2100  & 29.760  \\\\ $\\alpha$ & 90  & 90   & 90  & 90  & 90  & 90  \\\\ $\\beta$  & 90  & 110.15  & 99.38  & 95.74  & 99.58  & 90  \\\\ $\\gamma$ & 90  & 90   & 90  & 90  & 90  & 120  \\\\ ref.    & $^1$  & $^2$ & $^3$ & $^3$ & $^4$ & $^5$ \\\\ $N_{atoms}$ & 81& 80& 44& 88& 88& 66\\\\ \\hline \\end{tabular} \\\\ References: $^1$\\cite{CM02}, $^2$\\cite{CS73},$^3$\\cite{S76},$^4$\\cite{SF73},$^5$\\cite{B78} \\end{table*}   {\\it Ab initio} calculations of equilibrium isotope fractionation between minerals have received considerable attention recently. Previous studies, however, were mostly limited to simple crystals containing just a few atoms in the unit cell such as quartz, kaolinite or carbonate minerals \\citep{ML07,RC10}, as the methods used require considerable computational resources. Only very recently, the calculations have been extended to more complex crystalline solids containing up to 80 atoms in the unit cell by \\citet{SCh11}. There are different approaches used in the computation of the mass-dependent stable isotope equilibrium fractionation factors of minerals, but all methods require knowledge of the vibrational spectrum of the considered system, which is usually computed using {\\it ab initio} methods. \\citet{ML07} performed full normal mode analysis of the solid phases accounting for the phonon dispersion in reciprocal space. Because of the huge computational requirements, this method, although correct, can be only applied to the computation of stable isotope fractionation between simple phases. On the other hand, in order to derive the frequencies required for the computation of the fractionation factors \\citet{RC10} approximated the solids by small clusters and treated them as large molecules. This approach is based on well established theories of stable isotope fractionation \\citep{BM47,K82,CCH01} showing that the major contribution to the mass-dependent fractionation comes from the local vibrational motion of the fractionating element. In line with this finding \\citet{SCh11} has found that considering the phonon spectrum on a single phonon wave vector only is sufficient for modeling of $^{26}$Mg/$^{24}$Mg isotope fractionation between magnesium bearing crystal phases.  \\begin{figure*}[t] \\includegraphics[angle=270,width=4.in]{fig1.eps} \\caption{ The $\\beta$ factors for staurolite and spodumene. The lines represent the results using Eq. \\ref{BAPX} (solid lines) and Eq. \\ref{beta} with the full frequency spectrum (dotted lines). \\label{FIG1}} \\end{figure*}   As fluid-rock interactions are a major cause that alter the isotopic signature of a mineral in a rock, understanding the equilibrium isotope fractionation processes between minerals and aqueous fluids is of great importance in petrology. Although there has been considerable work on stable isotope fractionation between various minerals and the computational techniques are well established, the question of treating fluids, namely aqueous solutions remains open. Most of the {\\it ab initio} calculations of isotope fractionation in fluids use the cluster approach \\citep{DK08,HS08,Z09,RB10,RC10,HS10,YMMW01,Z10}, in which the considered species (ions or molecular complexes such as Fe, Mg, $\\rm H_3BO_3$) are surrounded by a hydration shell and the whole structure is relaxed assuming $T=0$~K. This approach is based on the computation of static atomic configurations and is valid at low temperatures only. In case of Li in aqueous solution at high temperatures ($T\\sim1000\\rm\\,K$), frequent exchange between particles of the hydration shell surrounding Li cation with the fluid is observed on time scales as short as picoseconds ($10^{-12}$~s, \\citet{JW09}). Distribution of cation coordination and cation-O bond lengths, effects that are expected to affect the isotope fractionation \\citep{BM47}, also change with pressure \\citep{JW09,WJ11}. These features are difficult to account for by using the cluster approximation for a compressible fluid at high temperature. The impact of the dynamical behavior of particles and compressibility of fluid must be investigated in order to properly compute the isotope fractionation in aqueous fluids. The only recent {\\it ab initio} work that accounts for the dynamical effects on the isotope fractionation in fluid is by \\citet{RB07} who considered boron isotope fractionation between $\\rm B(OH)_3$ and $\\rm B(OH)_4^-$ in aqueous solution. They performed {\\it ab initio} molecular dynamics simulations of this system and tried to use the vibrational density of states derived through the Fourier transform of the velocity auto-correlation function as an input for the calculation of the $^{11}$B/$^{10}$B isotope fractionation coefficient. The resulting fractionation factor $\\alpha=0.86$ is much lower than the experimental value $\\alpha=1.028$. Interestingly, the discrepancy between experiment and theory is cured by quenching the selected configurations along the molecular dynamics trajectory and computing the harmonic frequencies. The fractionation factor derived using these frequencies exactly reproduces the experimental value.    In this contribution we present an efficient approach to the computational prediction of equilibrium isotope fractionation between complex minerals and fluids at high $P$ and $T$. Both solids and fluids are treated as extended systems by application of periodic boundary conditions in all three spacial directions. We will demonstrate that at $T \\rm >600\\,K$ the fractionation factor can be computed by considering the force constants acting on the fractionating element only. Both solid and fluid supercells should be big enough to avoid significant interaction between atoms and their periodic images. In our investigation we use cells at least $5\\,\\rm\\AA$ wide in each spacial dimension. A representative distribution of relevant coordination environments in the fluid structure is obtained by performing Car-Parrinello molecular dynamics simulations \\citep{CPMD1}. For the calculation of the fluid fractionation factors, several random snapshots from this simulation are chosen. The force constants acting on the fractionating element and the resulting fractionation factors are then obtained for each configuration and the fractionation factor for the considered element in the fluid is computed as an average over the whole set of geometries.  As a test case for our approach, we have computed the fractionation factors between Li bearing aqueous fluids and three minerals, mica, staurolite and spodumene. For these systems, recent experimental data are available for comparison \\citep{W05,W07}. Furthermore, lithium as one of the lightest elements with two stable isotopes produces strong isotope signatures. It strongly fractionates into aqueous fluids during fluid-rock interaction processes and is used as a tracer of mass transfer in the subduction cycle \\citep{W05}. The two stable isotopes, $\\rm ^7Li$ and $\\rm ^6Li$, have respective abundances of 92.5$\\%$ and 7.5$\\%$. The large mass difference of $7.016003/6.015121=16.6\\%$ results in a prominent fractionation of at least a few {\\permil} even at high temperatures $T\\rm\\sim1000\\,K$. The experimental data on Li isotopes indicate a significant influence of the Li coordination and the Li-O bond length on the fractionation of Li isotopes. The heavier isotope preferentially occupies the lower coordinated sites and phases with shorter bond distance \\citep{WJ11}, which is expected also from the theoretical point of view \\citep{SMH09}. We will show that the application of {\\it ab initio} methods to Li-bearing silicates and fluids provides unique insight into the mechanisms driving equilibrium Li-isotope fractionation on the atomic scale.  \\begin{table*} \\caption{The $\\beta$ factors for mica (columns 2-5) and fractionation factors between mica and spodumene (last column) computed at $T\\rm=650\\,K$ for various mica polytypes and Li substitution sites. All values are given in $\\permil$. The measured value is that of \\citet{W07}. } \\label{T2} \\centering \\begin{tabular}{l c c c c c c}     % \\hline\\hline Mineral & Li1 & Li2 & Li3 & Average & $\\rm \\Delta^7Li_{mc.-spd.}$ \\\\ \\hline 1M & 13.9 & 14.9  & - & 14.6 & +4.7$\\pm0.9$  \\\\ occupation & 0.3 & 0.97 & -  & &   \\\\ 2M1 & 13.9 & 13.6 & - & 13.7 & +3.8$\\pm0.9$  \\\\ occupation & 0.38 & 0.92 & -  & &   \\\\ 2M2 & 13.8 & 13.4 & - & 13.6 & +3.7$\\pm0.9$  \\\\ occupation & 0.37 & 0.95 & -  & &   \\\\ 3T & 12.2 & 14.9 & 12.1 & 13.6 & +3.7$\\pm0.9$  \\\\ occupation & 0.7 & 0.89 & 0.14  & &   \\\\ exp. & &  & & & +2.5$\\pm$1.0 \\\\  \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "We propose a computationally efficient approach for computation of the $\\beta$ and isotope fractionation factors for complex minerals and fluids at high temperatures and pressures. We demonstrated that in order to derive the reliable $\\beta$ factors for either minerals or fluids at high $T$ and $P$ it is sufficient to know the force constants acting on the substituted isotope. This reduces significantly the computational time and allows for computations of isotope fractionation in complex materials containing even hundreds of atoms. In case of fluids we show that the widely used technique of representing aqueous solution as an ion-hydration-shell cluster is not sufficient to reproduce the isotope fractionation in aqueous solutions at elevated temperatures and pressures, when the dynamical character of the hydration shell and the compression of the fluid have to be accounted for. This can be achieved by {\\it ab initio} molecular dynamics simulation technique, which allows for direct access to the dynamical distributions of water (fluid) molecules around the considered ion and proper consideration of compression effects. The relevant isotope fractionation factors can be computed on a set of uncorrelated snapshot configurations extracted from the molecular dynamics trajectory.  We show that in the case of Li in aqueous solution it is sufficient to compute the $\\beta$ factors from the molecular dynamics simulations performed with a simulation cell containing a small number of atoms, which further reduces the computational time needed to perform the task. A system containing a single Li, a charge compensating anion and $8\\rm\\,H_2O$ molecules was sufficient to obtain the accurate $\\beta$ factors within the uncertainties of the {\\it ab initio} method used in the calculations.  We verify our approach by computing the Li isotopes fractionation factors between Li-bearing minerals and aqueous solutions and their comparison with the experimental data. The computed Li fractionation factors between staurolite, spodumene, mica and aqueous solutions reproduce the experimental results on quantitative and qualitative levels. We show that {\\it ab initio} calculations are able to predict the correct sequence of isotopes fractionation between considered materials as observed in the experiment. The computed fractionation factors are within $1\\,\\permil$ in agreement with the measured values. We also found that the thermal expansion of the solids affects the isotope fractionation process and its inclusion improves the agreement with the experimental data.  Our study shows that {\\it ab initio} computer simulations represent a powerful tool for prediction and understanding of equilibrium stable isotope fractionation processes between various phases including aqueous solutions at high pressures and temperatures. We expect that with the increasing power of computers and performance of the computational software these methods will be extensively applied to complement analytical techniques and to interpret measured isotopic signatures."
    },
    "1107/1107.1711_arXiv.txt": {
        "abstract": "Numerical calculations of merging black hole binaries indicate that asymmetric emission of gravitational radiation can kick the merged black hole at up to thousands of km/s, and  a number of systems have been observed recently whose properties are consistent with an active galactic nucleus containing a supermassive black hole moving with substantial velocity with respect to its broader accretion disk.  We study here  the effect of an impulsive kick delivered to a black hole on the dynamical evolution of its accretion disk using a smoothed particle hydrodynamics code, focusing attention on the role played by the kick angle with respect to the orbital angular momentum vector of the pre-kicked disk.  We find that for more vertical kicks, for which the angle between the kick and the normal vector to the disk $\\theta\\lesssim 30^\\circ$, a gap remains present in the inner disk, in accordance with the prediction from an analytic collisionless Keplerian disk model, while for more oblique kicks with $\\theta\\gtrsim 45^\\circ$, matter rapidly accretes toward the black hole.  There is a systematic trend for higher potential luminosities for more oblique kick angles for a given black hole mass, disk mass and kick velocity, and we find large amplitude oscillations in time in the case of a kick oriented $60^\\circ$ from the vertical.\\\\ ",
        "introduction": "In the past few years, a combination of surprising numerical studies and astronomical observations have indicated that asymmetric momentum losses in gravitational radiation from the mergers of  binary black holes (BH) may produce ``kicks'' of up to several thousand km/s.  For the mergers of supermassive black holes (SMBH) at the centers of galaxies, the kicks may be large enough to eject the remnant BH out of the galactic  center if not out of the galaxy entirely.  Kicked BHs may have already been observed indirectly as active galactic nuclei (AGN) with different components at different redshifts; broad-line regions that are thought to originate near the BH itself will remain bound to the kicked BH and acquire its new line-of-sight velocity, while narrowline systems that are produced further away will become unbound and remain behind.  While kicks from the mergers of unequal-mass, non-spinning BHs have long been predicted by post-Newtonian calculations \\citep{Fitchett:1984qn}, the ability to evolve black holes in fully general relativistic simulations has considerably expanded our view.  Indeed, it is the merger of equal-mass, {\\em spinning} BH that produce the largest kicks calculated to date, with maximum values of up to $4000$km/s possible for configurations with carefully chosen alignments \\citep{Campanelli:2007ew,Campanelli:2007cga,Gonzalez:2007hi}.  While speeds this large represent only a small fraction of the likely merging BH parameter space, Monte Carlo simulations of merging BHs with arbitrary spins of dimensionless magnitude $S/M^2=0.97$ find that the mean kick for  BH systems with mass ratios uniformly distributed between $0$ and $1$ is 630km/s, with more than 20\\% of the kicks larger than 1000km/s \\citep{Lousto:2009ka}.   Such kicks have the potential to unbind the remnant from smaller galaxies, or displace the BH and any bound gas for larger galaxies.  In roughly half of all major mergers between comparable-mass SMBHs, the merged SMBH should remain displaced for 30 Myr outside the central torus of material that would power an AGN \\citep{Komossa:2008as}.  Somewhat more recently, several AGN have been observed to contain broadline emission systems that appear to have very different redshifts than the narrowline emission systems.   In the first of these, SDSSJ092712.65+294344.0, a blueshift of $2650$km/s is observed for the broadline systems relative to narrow lines \\citep{Komossa:2008qd}.  Although several different physical models have been proposed to account for the observations, including broadline emission from the smaller SMBH within a binary surrounded by a disk \\citep{Dotti:2008yb,Bogdanovic:2008uz}, a pair of SMBHs at the respective centers of interacting galaxies \\citep{Heckman:2008en}, and even spatial coincidence \\citep{Shields:2008kn}, the kicked disk model remains entirely plausible.  Since then, several more systems with similar velocity offsets between different AGN emission regions have been discovered, including SDSS J105041.35+345631.3 [3500km/s; \\citep{Shields:2009jf}], CXOC J100043.1+020637 [1200km/s; \\citep{Civano:2010es}], and E1821+643 [2100 km/s; \\citep{Robinson:2010ui}].  Given a number of recent results that suggest that binary SMBH spins should align with the angular momentum of the accretion disk prior to merger, these large kick velocities are somewhat unexpected. Indeed, aligned spins can produce kicks of several hundred km/s \\citep{Campanelli:2007cga} but not the ``superkicks'' with recoil velocities $\\gtrsim 1000$km/s. \\citet{Bogdanovic:2007hp} suggest  that the accretion torques in gas-rich (``wet'') mergers should suffice to align the SMBH spins with the accretion disk prior to merger. In \\citet{Dotti:2009vv}, high resolution simulations of SMBH binaries whose orbits counter-rotate with regard to a surrounding disk indicate that they should undergo an angular momentum flip long before merger. With typical spin-orbit misalignments of no more than $10^\\circ-30^\\circ$ depending on the parameters of the disk, in particular its temperature, they find that the recoil kick during mergers should have a median value of $70$km/s, with superkicks an exceedingly rare event \\citep{Dotti:2009vz}.  While some of the observed candidate recoil velocities are so large that they fall well out into the tail of the kick velocity probability distribution function, it is difficult to constrain exactly how many systems indicate potential recoils given the challenges in clearly distinguishing multiple velocity components within a single AGN.  Given this difficulty, we attempt here to study potential electromagnetic (EM) signatures that would originate from post-merger disks.  The qualitative details of pre-merger evolution have been studied by other groups, and a relatively coherent picture emerges.  In general, each SMBH may be surrounded by a circum-BH accretion disk extending out to a distance substantially smaller than the binary separation, since orbits near the outer edge are unstable to perturbations from the other SMBH.  A circumbinary disk may be present as well, extending inward to a distance of roughly a few times that of the binary separation, with a gap appearing between the circumbinary disks and the inner disks.  Previous hydrodynamical studies have shown that  the inner edge of the circumbinary disk is driven into a high-eccentricity configuration that precesses slowly, while a 2-armed spiral density wave is formed extending out to larger radii \\citep{Macfadyen:2006jx}. Meanwhile, the inner disks will be fed by mass transfer from the circumbinary disk \\citep{Hayasaki:2006fq,Kimitake:2007fs}, as the $m=2$ azimuthal gravitational perturbation induces an elongation in the outer disk.  For circular orbits, the mass transfer rate is relatively constant, while for elliptic binaries the mass transfer takes on a more periodic character.  Finally, once the binary begins its gravitational radiation-driven plunge, the binary decouples form the outer disk, and mass transfer basically ceases.  Among the first predictions of the observational consequences of a post-kick accretion disk are those in \\citet{Loeb:2007wz}, who find that off-center quasars could be observed for up to $10^7$ years after a SMBH merger. \\citet{Schnittman:2008ez} find that the inner edge of the circumbinary disk is likely to occur at roughly 1000$M$, in typical relativistic units where $G=c=1$, with the value only weakly dependent on the typical disk parameters like the assumed $\\alpha$-parameter.  Based on semi-analytic models of the post-merger accretion disk, they find the potential for observable infrared emission lasting hundreds of thousands of years, leading to the prediction that several such sources might be present today, though they would be difficult to disentangle from other AGN sources.    On much shorter timescales immediately prior to the merger, the dissipation of gravitational radiation energy through spacetime metric-induced shearing of the disk can also lead to enhanced emission in the optically thin components of the disk \\citep{Kocsis:2008va}.  One of the first studies of post-kick disk dynamics was performed in \\citet{Shields:2008va}, who analyzed the approximate physical scales characterizing the post-merger disk and concluded that the total energy available to be dissipated by shocks is roughly $\\frac{1}{2}M_b v_{\\rm kick}^2$, where $M_b$ is the mass of the portion of the disk that remained bound and $v_{\\rm kick}$ is the SMBH kick velocity.  They estimated from a simple analytical model that an excess luminosity of $\\sim 10^{46}$ erg/s would be observable for roughly 3000 years, with a characteristic observed temperature of $\\sim 10^7$K assuming $v_{\\rm kick}=1000$km/s.  This model is checked by means of a {\\em collisionless} N-body simulation of a disk around a kicked SMBH, which found rough agreement with the analytical model and confirmed the prediction of a visible soft X-ray flare that would last for a few thousand years after the initial kick.  Using a 2-d version of the FLASH code, \\citet{Corrales:2009nv} are  able to study the response of a thin disk if the kick is  directed in the equatorial plane.  They find  that the characteristic response is  a one-armed spiral shock wave, capable of producing total luminosities up to $\\sim 10\\%$ of the Eddington luminosity on a timescale of months to years.  The relativistic decrease of the total SMBH mass, attributable to the energy carried off in gravitational waves and roughly $5\\%$ of the original total, leads  to a decrease in the luminosity of approximately $15\\%$ but does  not provide  a clear signature that can  be disentangled from the other global processes occurring in the disk.  Nevertheless, with future X-ray instruments such as the Square Kilometer Array, impulsive changes to a disk may be observable in the jet emission  \\citep{O'Neill:2008dg}. Even in the case where the BH kick is  very small, the secular, as opposed to dynamical, filling of the inner region of the disk should produce an afterglow that could exceed the Eddington luminosity if the accretion rate is  sufficiently large prior to the binary decoupling and merger \\citep{Shapiro:2009uy}.  The most direct comparison to the calculations we present here is found in \\citet{Rossi:2009nk}.  Using an analytic treatment of disks with power-law density profiles, they construct a model for disk evolutions in which, immediately after the kick, fluid elements are assumed to circularize at radii determined by their specific angular momentum, with the resulting energy gain by shocks released as EM radiation.  Their work establishes that the primary energy reservoir for kicked disks  is potential energy that can be released by elliptical orbits, not the relative kinetic energy of the kick itself nor the impulse sent through the disk by instantaneous mass loss to gravitational radiation by the central SMBH during the merger.  They also perform detailed 2-d Eulerian and 3-d SPH simulations of post-kick accretion disks, although there are some important differences between the latter simulations and most of our SPH studies, as we discuss throughout the paper below.   Among these, they assume that disks are isothermal, with all shock heating immediately radiated away, whereas we evolve the energy equation and allow the fluid to heat.  They also include an ``accretion radius'', $R_{acc}$, such that particles that approach closer than that distance to the black hole are accreted and removed from the simulation domain (G. Lodato, private communication), whereas all particles remain in our simulation throughout the duration of a run.  As such, their results and ours  bracket the range of possible heating scenarios.  The 2-d calculations by \\citet{Rossi:2009nk} of razor-thin disks using the ZEUS code indicate that vertical kicks (that is, in the direction of the disk angular momentum vector) lead to modulated emission, unlike all other kick angles they consider. This result is confirmed by their 3-d SPH simulations. In-plane kicks develop a clear spiral-wave structure with accretion streams forming as the simulation progressed, but a smooth luminosity profile.  Their 3-d SPH simulations with oblique (less vertical) kicks  indicate essentially a 2-phase model for the observed EM emission.  Immediately after the kick, the majority of the luminosity may be attributed to the innermost region of the disk dissipating the kinetic energy it acquired during the kick, while at later times, after roughly one standard timescale of the disk ($\\hat{t}=1$ in the notation of Eq.~\\ref{eq:tscale} below) and thereafter, the dominant dissipation mode is potential energy from infalling material on elliptical orbits.  Post-kick disks are found to be rather compact, extending out to roughly the ``bound radius'' ($\\hat{r}=1$ according to the notation of our Eq.~\\ref{eq:rscale}) with steep density dropoffs at larger radii.  Luminosities are generally highest for in-plane kicks, with roughly a factor of four difference in peak luminosity between largest peak luminosity (in-plane kick) and the smallest (vertical kick), with peaks occurring later for more oblique kicks.  Relativistic BH mass decrease was found to be unimportant at large radii, and potentially important only in the vicinity of the BH (out to a few hundred Schwarzschild radii) where the effect of the kick is merely perturbative compared to the nearly relativistic Keplerian velocities.  Numerical relativity groups have also considered the hydrodynamics of matter around both binary SMBHs and kicked BHs, typically at much smaller size scales.  In \\citet{Bode:2009mt}, flows of gas  around a binary SMBH system are  considered at scales roughly $10^5$ smaller than typical Newtonian calculations, spanning scales roughly 1AU across, rather than $\\sim$parsec scales.  They find  that EM emission is  dominated by variability created by Doppler beaming of the SMBHs as they shock the gas, leading to an EM signal that demonstrates the same periodicity as the gravitational wave signal, with corresponding peaks in the timing of the maximum emission for each.  In a follow-up work, Bode and collaborators \\citep{Bode:2011tq} predict that observable EM emission from near the SMBHs is much more likely to arise in a hot accretion flow, in which a flare would be seen coincident with the merger. In \\citet{Megevand:2009yx},  the effect of a kicked BH moving through the equatorial plane of an accretion torus is  considered using a fully general relativistic Eulerian hydrodynamics code.  In their simulations, the newly-merged SMBH is  surrounded by a torus extending out to 50M (50AU for a $10^8M_\\odot$ SMBH),  and the overall timescale studied is  approximately 10,000M (in relativistic units with $G=c=1$), or about two months.  They find  that a kick in the direction of the equatorial plane of the torus produces the strongest shock in the system and therefore the strongest EM emission, consistent with studies that examine disks on substantially larger scales.  By ray-tracing their simulations in a post-processing step \\citep{Anderson:2009fa}, they confirm that simple Bremsstrahlung luminosity estimates yield a qualitatively accurate picture of the disk luminosity for high-energy radiation, while their torus is optically thick to low-energy emission.  Numerical calculations of vacuum EM fields surrounding an SMBH merger indicate that they could contribute to periodic emission \\citep{Palenzuela:2009yr,Palenzuela:2009hx} but are likely to be too small in amplitude and at the wrong frequencies ($\\sim 10^{-4}$Hz) to be observed directly \\citep{Mosta:2009rr}.  Such mergers could produce observable levels of Poynting flux in jets \\citep{Palenzuela:2010nf}, however, through a binary analogue of the Blandford-Znajek mechanism \\citep{Blandford:1977ds}, which seems especially effective for spinning BHs \\citep{Neilsen:2010ax}. Calculations of non-Keplerian accretion disks in general relativity indicate  that the spiral wave structures seen in Newtonian simulations could exist in relativistic models with small disks when  the BH kick is  sufficiently small, but that larger kicks disrupt the spiral pattern, as could dissipative processes such as magnetic stress or radiative cooling \\citep{Zanotti:2010xs}. The inferred emission due to synchrotron emission from a relativistic disk is  considered by the Illinois group \\citep{Farris:2009mt,Farris:2011vx}, who find  that emission could peak at a luminosity of $\\sim 10^{46}$ erg/s, a few orders of magnitude brighter than the corresponding bremsstrahlung luminosity and potentially observable by either WFIRST or LSST.  The outline of the paper is as follows.  In Section~\\ref{sec:scales}, we introduce the physical scales that define the kicked disk problem, and discuss  the basic dynamics of disks around a kicked black hole when hydrodynamic effects are ignored.  In Sec.~\\ref{sec:SPH}, we describe the SPH code used to perform 3-dimensional  simulations, including versions that seek to replicate previous results and those that use new techniques to bracket the range of physical physical predictions.   Results from these simulations, focusing on the hydrodynamic and thermodynamic evolution of the disk, are reported in Sec.~\\ref{sec:runs}.  Finally, in Sec.~\\ref{sec:discussion}, we lay out consequences of our results and plans to extend them in the future. ",
        "conclusions": "\\label{sec:discussion}  In this paper, we have studied the response of a quasi-equilibrium accretion disk around a SMBH that undergoes an impulsive kick, presumably because of a binary merger and the corresponding asymmetric emission of linear momentum in gravitational waves.  While there are a number of sources that have been identified as candidate kicked disks with recoil velocities so large that they would be in the far tails of the kick velocity distribution, our results here are scale-free with regard to the kick, and may be applied to a broad swath of potential kicked systems.  Indeed, while for our assumed reference model with $M_{\\rm BH}=10^8M_\\odot$, $m_{disk}=10^4M_\\odot$, and $v_{\\rm kick}=1000$ km/s we find characteristic timescales of roughly 1000 years and disk luminosities and temperatures of up to $10^{42}$erg/s and $10^{8-9}$K, respectively, the results should work for a wide range of masses and kick velocities. There is a trade-off, to be had, of course, especially given the dependence of many of the quantities we investigate on the kick velocity.  Indeed, cutting the kick velocity by half and leaving the masses of the BH and disk unchanged increases the timescales we consider by a factor of eight while cutting the energies and temperatures by a factor of four and thus the luminosities by 32, indicating that there should be a definite observational bias toward the larger kicks.  The code we introduce uses 3-dimensional SPH techniques, and is modified to incorporate a Lagrangian (``grad-h'') prescription that allows energy to be conserved to high precision over the course of all our runs.  We also introduce a Lagrangian-based black hole smoothing potential, which proves critical in allowing us to avoid numerical issues associated with point-like potentials.   One of the most important purely numerical conclusions of this work is that black holes must be handled extremely carefully in SPH, since they can introduce particle clustering instabilities that are both conservative but highly unphysical.  This treatment, along with the inclusion of shock-heating, has allowed us to perform what we believe to be the first 3-d SPH simulations of these systems that do not require imposing an ad hoc ``accretion radius'' inside of which SPH particles are removed from the simulation.  In order to examine the phase space of post-kick disks, we have varied the kick angle of the SMBH with respect to the initial disk orientation, which has a large effect on the resulting evolution.  More oblique kicks, i.e., those most oriented toward the equatorial plane of the original disk, produce substantially higher peak luminosities for a given BH and disk mass and kick velocity, with roughly a factor of four gain between kicks oriented $15^\\circ$ away from vertical and those oriented at $60^\\circ$ from vertical.  Assuming an astrophysical context in which the disk is aligned with the SMBH binary orbit prior to merger, it is unlikely that more oblique kicks will actually yield more luminous events, however.  Kick velocities are systematically higher for more vertical kicks \\citep{Lousto:2009ka}, and given the $v_{kick}^5$ dependence of the luminosity, the brightest kicked disks are likely to be those with the highest speed kicks, with the kick angle playing a secondary role.  We find more rapid luminosity peaks appearing for more oblique kicks, which is the opposite result from a previous set of simulations that considered power-law density profile disks \\citep{Rossi:2009nk}. Based on our $60^\\circ$ calculation, we attribute this to the rapid filling of the innermost region of the disk by material kicked into high eccentricity orbits with extremely small periastron distances (see Fig.~\\ref{fig:rperi}), which flows inward from the inner edge of the pre-merger disk; we note that this pattern differs from that found in Fig.~15 of \\citet{Rossi:2009nk}, likely due to the combined effects of shock heating and their excision of particles from the innermost regions of the disk.   Clearly, one of the important conclusions of this work and others is that flows of material toward the SMBH after the kick can release tremendous amounts of energy, but need to be modeled very carefully to derive reasonable light curves and spectra, pushing the limits of current numerical simulations.  The surface density oscillations and the resulting quasi-periodic emission we observe in time for the most oblique kicks will make an interesting topic for further research.  Based on our ``gap-filling'' model  it seems clear the dynamics of the inner disk play an important role in the dynamical evolution and EM emission from the disk, and it would be highly useful to extend the same simulations on both sides of the kick to explore the full history of a circumbinary/post-kick disk. In particular, it will be important to accurately resolve the inner edge of the circumbinary disk, where the tidal field of the binary should have swept out a gap \\citep{Schnittman:2008ez}, and to understand how infalling material from that region interacts with the more tenuous tidal streams of matter in the process of accreting onto the SMBH prior to their merger \\citep{Macfadyen:2006jx,Hayasaki:2006fq,Kimitake:2007fs}. It will also be useful to incorporate a more thorough treatment of the radiative evolution of the disk, since a proper treatment of radiative cooling and disk opacities will break the scale-invariance of the results and allow for much more accurate predictions of observable phenomena."
    },
    "1107/1107.4302_arXiv.txt": {
        "abstract": "{ NGC 6251 is a luminous radio galaxy $\\approx 104$ Mpc away that was detected significantly with the \\emph{\\it Fermi Gamma-ray Space Telescope}, and before that with EGRET (onboard the {\\it Compton Gamma-ray Observatory}). Different observational constraints favor a nuclear origin for the $\\gamma$-ray emission. Here we present a study of the spectral energy distribution (SED) of the core of NGC 6251, and give results of modeling in the one-zone synchrotron/SSC framework. The SSC model provides a good description of the radio to $\\gamma$-ray emission but, as for other misaligned sources, predicts a lower Lorentz factor ($\\Gamma\\sim2.4$) than typically found when modeling blazars. If the blazar unification scenario is correct, this seems to point to  the presence of at least two emitting regions in these objects, one with a higher and one with a lower Lorentz factor. The solution of a structured jet, with a fast moving spine surrounded by a slow layer, is explored and the consequences of the two models for the jet energetics and evolution are discussed.} ",
        "introduction": "The  \\emph{Fermi Gamma-ray Space Telescope} has confirmed misaligned AGNs (MAGNs), including radio galaxies (RGs) and steep--spectrum radio quasars (SSRQs), as a new and important class of $\\gamma$--ray emitters \\citep{2010ApJ...715..429A,2010ApJ...720..912A}. In the first year of activity, the Large Area Telescope \\citep[LAT,][]{2009ApJ...697.1071A} onboard \\fermi detected 11 MAGNs belonging to the 3CRR, 3CR, and  MS4 catalogs at 178 and 408 MHz. At these low frequencies, the Cambridge and Molonglo surveys favor the detection of radio galaxies over  blazars. Seven of the 11 \\fermi MAGNs are nearby ($z\\la 0.06$) Fanaroff-Riley type I \\citep[FRI;][]{1974MNRAS.167P..31F} radio galaxies, while the rest are FRII radio galaxies.\\\\ \\noindent The FRI radio galaxies detected with \\fermi are significantly  less $\\gamma$-ray luminous than  BL Lac objects. The BL Lac objects have isotropic 100 MeV -- 10 GeV luminosities $L_\\gamma \\approx 10^{44} - 10^{46}$ erg s$^{-1}$, significantly higher than the FRI radio galaxies, which have $L_\\gamma \\approx 10^{41} - 10^{44}$ erg s$^{-1}$. They also have spectral slopes that are consistent with low- or intermediate-synchrotron-peaked blazars \\citep[][]{2010ApJ...715..429A,2010ApJ...716...30A}. This is consistent with AGN unified schemes, according to which an increase in the inclination angle of the jet axis with respect to the observer's line-of-sight implies a de-amplification of the observed flux and a general shift in the emission to lower energies \\citep[][and references therein]{1984ApJ...280..569U,1995PASP..107..803U}. The classification itself of misaligned sources is relative to our expectations of detecting in radio galaxies strongly beamed jet radiation from highly relativistic ($\\Gamma\\gtrsim10$) bulk flows, as in the case of blazars. Thus, for the sake of clarity, by misaligned here we do not mean that the observer line-of-sight is outside the radiation cone of the source but rather that it falls outside the (narrow) beaming cone we would expect in the case of a highly relativistic core.\\\\ A full comparison with the expectations of the unification scenarios for AGNs has not yet been considered, but the addition of a substantial number of misaligned AGNs at 100 MeV -- 10 GeV energies now enables this to be made.  \\noindent Here we present our study of the SED of the core of NGC 6251, identified with 1FGLJ1635.4+8228 in the First \\fermi-LAT Catalog \\citep[][First \\fermi-LAT Source Catalog, 1FGL, and First \\fermi-LAT AGN Catalog, 1LAC, respectively]{2010ApJS..188..405A,2010ApJ...715..429A}. NGC 6251 is the fifth $\\gamma$-ray brightest radio galaxy in the MAGN sample with an integrated flux $F_{-9}=(36\\pm8)$ in units of $10^{-9}$ ph($> 100$ MeV) cm$^{-2}$ s$^{-1}$. It is associated with the EGRET source 3EG J1621+8203 \\citep{2002ApJ...574..693M}, being after Centaurus A one of the most likely associations of an EGRET source with a radio galaxy (also of interest are NGC 1275, 3C 111, and M87). Being bright and relatively close \\citep[$z=0.02471$][]{2003AJ....126.2268W}, NGC 6251 provides a nearby cosmic laboratory to explore the jet structure of radio sources hosting a supermassive black hole.\\\\ After describing the source in more detail (Sect. 2), we describe the multiwavelength datasets used to assemble the spectrum of NGC 6251 (Sect. 3). In Sect. 4, we illustrate some uncertainties regarding jet orientation and location, before presenting the results of one-zone synchrotron self-Compton (SSC) modeling in Sect. 5. There we derive magnetic fields, outflow Lorentz factors, and absolute power estimates for comparison with blazars. Parameter comparisons between SSC models of  BL Lacs and FRI galaxies reveal the incongruity between these two classes of sources, which we discuss in light of the unification scenario and explore the feasibility of an alternative scenario. In Sect. 6, we consider the connection between the jet power, the disk luminosity, and the accretion. The results are summarized in Sect. 7. ",
        "conclusions": "We have presented a study of the broad-band nuclear SED of the radio galaxy NGC~6251, one of the MAGNs detected by {\\it Fermi}-LAT. In agreement with previous studies, the nuclear SED is dominated by non-thermal emission related to the sub-pc jet. A nuclear origin appears to be the most likely hypothesis for the $\\gamma$-ray emission. This is also supported by the results of our X-ray analysis of archival XMM-Newton, Chandra, and Swift observations. Both models, SSC and spine-layer, adopted to reproduce the observed SED, provide a good overall fit. In the following, we summarize the main results of the SED modeling.\\\\ \\\\ A single-zone SSC model assumes a slow and heavy jet. The particle power dominates over the magnetic field power by about three orders of magnitude. This seems to rule out a magnetically accelerated and confined jet. The X-ray halo surrounding the jet can account for its confinement on kpc scales but the core region remains poorly constrained. A low bulk Lorentz factor, when a one-zone SSC model is adopted, is also shared with other MAGNs detected by {\\it Fermi}-LAT \\citep[][]{2009ApJ...699...31A,2009ApJ...707...55A,2010ApJ...715..554K,2010ApJ...719.1433A}. This suggests the presence of at least a second site of $\\gamma$-ray emission in addition to the one we observe when the jet is pointing directly toward us \\citep{2000A&A...358..104C}. A significant reduction of the jet inclination angle, which would allow higher bulk Lorentz factors, is unsupported by the radio observations. These results should not be significantly affected by an overestimate of the $\\gamma$-ray flux of a factor of a few ($\\approx 5$). \\\\ A structured jet with a fast inner component and a slower layer is a possible explanation. In the spine-layer model proposed in \\citet{2005A&A...432..401G}, the jet is relatively light and could be magnetically confined. The fast inner component  carries the bulk of the jet power. On the other hand,  the strong radiative dissipation ($L_r\\sim0.3\\,L_{kin}$) is at odds with the Mpc-length of the jet and the flux variability predicted by the model in the high energy band has not yet been observed.\\\\ The derived jet powers, $\\approx 10^{44}$ -- $10^{45}$ erg s$^{-1}$, are model-dependent and rely on some important assumptions. However, similar values were derived using the 151 MHz luminosity \\citep{1999MNRAS.309.1017W}.\\\\ The disk component appears to be completely hidden by the jet emission. The sub-Eddington luminosity regime ($L_{disk}/L_{Edd}\\lesssim 10^{-4}$) could be related to the presence of a RIAF. Support to this hypothesis is provided by the estimates of the Bondi accretion rate, albeit with significant uncertainties in the ISM physical parameters. As has been found for blazar sources, the mechanism channeling the accretion power into the jet seems to work in a more efficient way ($L_{kin}/L_{disk}\\gtrsim10$ and $L_{kin}\\sim0.1-1 P_{accr}$)."
    },
    "1107/1107.4244_arXiv.txt": {
        "abstract": "Theoretical predictions for the ensemble quasar structure function are tested using multi-epoch observations of Stripe 82 collected by the Sloan Digital Sky Survey.  We reanalyze the entire available volume of the $g$-band imaging data using difference image photometry and build high quality light curves for 7562 spectroscopically confirmed quasars.  Our structure function includes $\\sim4.8\\times10^6$ pairs of measurements and covers a wide range of time lags between 3 days and 6.9 years in the quasar rest frame. A broken power-law fit to this data shows the presence of two slopes $\\alpha_1=0.33$ and $\\alpha_2=0.79$ with the break at $\\sim42$ days.  The structure function compiled using only flux increases is slightly lower than that for variations of the opposite sign, revealing a slight asymmetry between the leading and trailing edge of a typical flare. The reality of these features is confirmed with monte-carlo simulations.  We give simple interpretation of the results in the frames of existing theoretical models. ",
        "introduction": "\\label{sec:intro}  Several decades after the discovery of quasars, the physical origin of their variability is still not fully understood. An unambiguous explanation of the intrinsic variations in quasar emission will likely provide the key to understanding the energy supply of quasars. The most frequently invoked model connects variability of quasars to instabilities arising in the accretion disk around a super-massive black hole (e.g.~\\citealt{1998ApJ...504..671K,2006ApJ...642...87P}). In the starburst model quasar light curves are a superposition of very frequent supernova explosions in the host galaxy (e.g.~\\citealt{1992MNRAS.255..713T,1996MNRAS.282.1191C,1997MNRAS.286..271A}). Gravitational microlensing by compact bodies within the host---such as normal stars or certain candidates for dark matter---also produces observable flux changes (e.g.~\\citealt{2002MNRAS.329...76H,2007A&A...462..581H}), but their contribution to the observed level of variability remains uncertain.  The models can be distinguished based on their predictions for the dependence of the slope and the amplitude of the structure function on the time-scale of variability in the quasar rest frame.  Recently, individual quasar light curves have been successfully modeled as a damped random walk (DRW) process (\\citealt{2009ApJ...698..895K}, \\citealt{2010ApJ...708..927K}, \\citealt{2010ApJ...721.1014M}).  While the method allows one to study the structure function of individual objects as a function of black hole mass, quasar luminosity, wavelength of observation (\\citealt{2010ApJ...721.1014M}), it also fixes the slope of the power spectrum in the high frequency regime at $-2$ for every quasar.  Moreover, the DRW model excludes the possibility that the structure function of positive flux variations differs from the one for negative changes. In this paper we rely on more traditional statistics using ensemble structure functions.  A recent influx of survey data on variability of tens of thousands of QSOs is fueling numerous observational studies and calls for a new look at the proposed scenarios.  The most recent papers favor the disk instability model (\\citealt{2009ApJ...696.1241B}, \\citealt{2005AJ....129..615D}, \\citealt{2004ApJ...601..692V}) with the exception of \\cite{2002MNRAS.329...76H} who find evidence in support of the microlensing model.  Here, we re-analyze the repeated $g$-band imaging of Stripe 82 collected by the SDSS using difference image photometry and build an ensemble structure function over a wide range of time-scales in order to test the theoretical predictions.  Our results show that in their present form none of the considered models explains the behavior of the structure function on all time-scales. Instead, a hybrid model including both a starburst and disk instabilities naturally follows from the data.  The paper is organized as follows. In Section~\\ref{sec:data} we describe the data and a method of extracting light curves. Basic properties of the structure function are given in Section~\\ref{sec:sf}.  In Section~\\ref{sec:sf_ensemble} we build the structure function for the ensemble of quasars.  We discuss the results in Section~\\ref{sec:discussion} and a summary is given in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  We consider three processes frequently invoked to explain the variability of quasars: instabilities in the accretion flow around a supermassive black hole, supernova explosions related to the starburst phenomenon, and gravitational microlensing by compact bodies in the host galaxy \\citep{1998ApJ...504..671K,2002MNRAS.329...76H}. SFs of light curves predicted by each model are characterized by different slopes. The expected SF slopes are $0.25\\pm0.03$, $0.44\\pm0.03$, $0.83\\pm0.08$ respectively for variability generated by miscrolensing, disk instability, and starburst mechanisms \\citep{2002MNRAS.329...76H}.  The SF slopes found in Section~\\ref{sec:sf_ensemble} are 0.79 and 0.33 correspondingly for time lags below and above $42$ days. Taken at face value, the first slope is consistent with the starburst model, and the second slope falls half way between the predictions of the disk instability model and the microlensing model.  Note that SF slopes in the range 0.30--0.35 have been obtained in other works, but the authors still attributed the variability to disk instabilities.  \\cite{2009ApJ...696.1241B} and \\cite{2011A&A...525A..37M} present an overview of recent results.  \\begin{figure} \\includegraphics[width=\\columnwidth]{f8_sf_plus_minus.eps} \\figcaption{Testing the asymmetry in the quasar structure function between the leading and trailing branch of a typical flare. Flux increases contribute to $S_+$ (blue) and flux drops contribute to $S_-$ (green). \\label{fig:sf_plus_minus}} \\end{figure}  Asymmetries in the SF provide an additional test of model predictions.  The SF that only includes those pairs of measurements for which the flux increases with time, $S_+$, may be different from $S_-$, the SF characterizing a decrease in flux. \\cite{1998ApJ...504..671K} show that in starburst models $S_+>S_-$. Disk instabilities produce $S_->S_+$ and microlensing is symmetric, i.e. $S_-=S_+$ when averaged over sufficiently long time intervals \\citep{2002MNRAS.329...76H}. The difference between $S_-$ and $S_+$ depends on the parameters of the model and therefore $S_- \\simeq S_+$ is still possible in both the starburst and the disk instability model \\citep{1998ApJ...504..671K}.  Noise-corrected $S_+$ and $S_-$ functions are shown in Figure~\\ref{fig:sf_plus_minus} with blue and green lines respectively. For long time lags in the range 300--1600 days, where the SF slope is $\\sim0.33$, we have $S_- > S_+$, as predicted by the disk instability model. The significance of the difference is not very high, but allows us to exclude the microlensing model. The difference disappears for time-scales below 100 days, where the slope is consistent with the starburst model. This may suggest that a typical supernova rate in a starbursts in the host galaxies of quasars is as high as $\\sim100$ yr$^{-1}$ and effectively washes out any detectable asymmetry (see Fig. 6 in~\\cite{1998ApJ...504..671K}). However, the rate should be treated with caution since it is not derived from direct measurements. Moreover, the predictions for $S_+$ and $S_-$ were derived from simulations of a single variability mechanism. The addition of another variability process can wash out the difference between $S_+$ and $S_-$.  This is an interesting topic for future research.  The difference in $S_+$ and $S_-$ was also investigated by \\cite{2005AJ....129..615D} and \\cite{2009ApJ...696.1241B}. The former study found tentative evidence that $S_+ > S_-$ (light curves have fast rise and slow decay as in starburst model) for 400 days $\\le\\tau\\le$ 1500 days, while in the latter work $S_+ \\approx S_-$ over the full range of measurements ($\\tau\\le1000$ days).  The explanation of the difference in results is difficult because each paper is based on entirely different data and requires further research. The \\cite{2005AJ....129..615D} analysis relies on cross survey flux comparisons.  The introduction of the supernova variability mechanism may be redundant as the shape of the SF can be explained by detailed variability mechanisms acting in accretion disks. For instance, the PSD of X-ray light curves of some active galactic nuclei have two (e.g. ~\\cite{2007MNRAS.378..649S}) or more distinct slopes~\\citep{2007MNRAS.382..985M}. Considering that the optical emission tends to follow the X-ray flux with some time delay one may expect that PSDs of optical light curves will display two slopes. Moreover, the SF shape found here is in qualitative agreement with the model proposed in~\\cite{2009MNRAS.397.2004A}, where large amplitude optical variations on the time scales of hundreds of days are attributed to the fluctuations in accretion rate and short-term (order of days), small amplitude variations arise due to reprocessing of X-rays.        We applied the image differencing technique to the entire $g$-band imaging data set of SDSS Stripe 82 and constructed high quality light curves for 7562 spectroscopically identified quasars with no less then 20 epochs. The variability analysis was performed using the ensemble structure function. While the shape and normalization of the quasar structure function derived here are in perfect agreement with the results of \\cite{2008MNRAS.383.1232W} based on an earlier SDSS data release, our structure function is estimated from a much larger data set and covers a substantially wider range of time lags from 8 hours to 6.9 years in the quasar rest frame.  We find that the ensemble SF reveals two distinct power-law slopes with the break around $42$ days and confirm the presence of these features with monte-carlo simulations. The presence of two slopes in the ensemble SF is independent on the black hole mass range of quasars used to build it. The slopes estimated from a broken pawer-law fit to the data are $\\alpha_1\\simeq0.79$ for $\\tau\\le100$ days and $\\alpha_2\\simeq0.33$ for $\\tau\\ge300$ days.  Using predictions from the theoretical models, the variability of quasars on times scales $\\tau\\le100$ days can be explained by a starburst model. The slope of the structure function on time-scales longer than about 3 months is rather close to models where variability is explained by gravitational microlensing.  However, on time-scales $\\tau\\ge300$ days we detect a significant asymmetry in the SF characteristic for disk instabilities that rules out a substantial contribution of microlensing.  It is also possible to explain the shape of SF without invoking the starburst model, but using the model, where short term optical variations are caused by X-rays reprocessing and long therm variations originate from optical emitting regions of accretion disk~\\citep{2009MNRAS.397.2004A}.   The existing SDSS Stripe 82 data offer a great opportunity to better understand the origin of variability in quasars. The main limiting factor at present is the lack of detailed theoretical predictions that go beyond the basic characteristics such as the SF slope and include the detailed shape and normalization of the SF, as well as asymmetries and color information."
    },
    "1107/1107.3652_arXiv.txt": {
        "abstract": "We show that the distribution of luminosities of Brightest Cluster Galaxies in an SDSS-based group catalog suggests that BCG luminosities are just the statistical extremes of the group galaxy luminosity function.  This latter happens to be very well approximated by the all-galaxy luminosity function (restricted to $M_r<-19.9$), provided one uses a parametrization of this function that is accurate at the bright end. A similar analysis of the luminosity distribution of the Brightest Satellite Galaxies suggests that they are best thought of as being the second brightest pick from the same luminosity distribution of which BCGs are the brightest.  I.e., BSGs are not the brightest of some universal satellite luminosity function, in contrast to what Halo Model analyses of the luminosity dependence of clustering suggest. However, we then use mark correlations to provide a novel test of these order statistics, showing that the hypothesis of a universal luminosity function (i.e. no halo mass dependence) from which the BCGs and BSGs are drawn is incompatible with the data, despite the fact that there was no hint of this in the BCG and BSG luminosity distributions themselves. We also discuss why, since extreme value statistics are explicitly a function of the number of draws, the consistency of BCG luminosities with extreme value statistics is most clearly seen if one is careful to perform the test at fixed group richness \\Ngal.  Tests at, e.g., fixed total group luminosity $L_{\\rm tot}$, will generally be biased and may lead to erroneous conclusions. ",
        "introduction": "The study of whether or not the brightest cluster galaxy (the BCG) is special, rather than simply being the brightest by chance, has a long history (Scott 1957; Schechter \\& Peebles 1976; Tremaine \\& Richstone 1977; Bhavsar \\& Barrow 1985; Loh \\& Strauss 2006; Lin, Ostriker \\& Miller 2010, Dobos \\& Csabai 2011). Since the BCG is usually at or close to the cluster center, it is plausible that its formation history was dominated by different physical processes compared to the other galaxies in the cluster, which we will call satellites.  For this reason, models of BCG formation routinely assume that BCGs are special (e.g. Milosavljevi\\'c, \\etal\\ 2006; De Lucia \\& Blaizot 2007). And there is growing evidence from scaling relations that BCGs are indeed special:  they have larger sizes, smaller velocity dispersions, smaller color gradients, and are less spherical than expected for their luminosities (e.g. Bernardi \\etal\\ 2011).  But whether or not this has left a distinct signature in the distribution of BCG luminosities is an open question.  There are at least two reasons why the answer is interesting. First, Halo Model (see Cooray \\& Sheth 2002 for a review) interpretations of the luminosity dependence of galaxy clustering explicitly distinguish between the central and the other (satellite) galaxies in a halo. If BCGs are just the brightest of a universal galaxy luminosity function, then including this constraint significantly simplifies such analyses.  The second is that the distribution of BCG luminosities is considerably narrower than that of all galaxies, so the large number of clusters that will soon be available may permit BCGs to provide a useful consistency check of standard candle constraints on the luminosity-distance relation (Hubble 1936; Scott 1957; Sandage, Kristian \\& Westphal 1976). Alternatively, if the $d_{\\rm L}(z)$ relation is known and BCGs are just statistical extremes, then the BCG luminosity function can be used to provide a rather accurate determination of the evolution of the galaxy luminosity function.  Since BCGs are bright, they are amongst the easiest targets for spectroscopy, so this measure of galaxy evolution comes at a small fraction of the cost of a full galaxy survey. On the other hand, if they are special in other ways, one must also understand how the physical processes which affected their formation evolve.  In what follows we will use the Halo Model analyses to motivate why one might have thought that the BCG luminosity function might simply be given by the extreme value statistics of the all galaxy luminosity function, even though centrals are often explicitly treated as being special. In this case, it is natural to ask if the luminosity function of the brightest satellite galaxies -- the BSG luminosity function -- is simply given by the associated order statistics of the second brightest (rather than the brightest) object. Then we will show that Halo Model analyses also suggest that the BSG luminosity function should be given by the extreme value statistics of a satellite galaxy luminosity function, rather than the order statistics of the all galaxy luminosity function. We will show that both statements cannot be right, so we turn to the data to decide the issue.  Section~\\ref{prelims} summarizes the main results from Order Statistics and the Halo Model approach which are relevant to this study. Section~\\ref{sdss} uses a group catalog of SDSS galaxies from Berlind \\etal\\ (2006) (henceforth B06) to test these models regarding BCGs and BSGs. These tests include the luminosity functions of BCGs and BSGs, the ratio of the luminosities of the first and second brightest galaxies in a cluster, and the luminosity function of satellite galaxies (i.e., not just the BSGs). Section~\\ref{marks} shows that luminosity weighted clustering provides a novel test of Order Statistics. A final section summarizes. Appendix~\\ref{psatExstat} provides a derivation of the satellite distribution if the BCGs satisfy extreme value statisics. Appendix~\\ref{shuffled} demonstrates that tests of order statistics are better made holding the number of galaxies fixed, rather than other parameters such as group luminosity, etc. And Appendix~\\ref{zehaviHOD} describes the details of the Halo Model that are relevant to our study. ",
        "conclusions": "Brightest cluster galaxies are interesting for several reasons, the simplest being that they are the brightest and therefore most easily observable of galaxies. As mentioned in the Introduction, it is also becoming increasingly certain that they are physically distinct in several other ways, e.g. in their morphology, velocity dispersions, etc., from the other galaxies (the satellites) which reside in groups and clusters. We argued that, from the point of view of Halo Model analyses, it is interesting to ask whether the luminosities of the BCGs retain any signature of their special nature. Our analysis shows that, at the level of one-point distributions, the answer appears to be `no': BCG luminosities (insofar as they are accurately represented in the Mr20 catalog that we analysed) are consistent with being the statistical extremes of a universal luminosity function, which is well-described by \\eqn{pgal-bernardi}.  We also analysed the luminosity functions of the satellites in the Mr20 catalog. The behaviour of these distributions for different values of group richness (i.e. number of members) $N$ is especially interesting, since it potentially allows us to distinguish between two seemingly reasonable but incompatible predictions of Halo Model analyses. On the one hand, since the BCG luminosities are simply statistical extremes of a universal galaxy luminosity distribution, it is reasonable to expect that (a) the satellite luminosity function is just that given by subtracting the BCGs from the universal one (equation~\\ref{app-psat}), and (b) the \\emph{brightest} satellite (BSG) luminosities are simply the \\emph{second} brightest of $N$ draws of the universal galaxy distribution. On the other hand, the HOD model predicts that (c) the satellite function is itself approximately universal, and (d) the BSGs are statistical extremes of this function. We showed that these two sets of predictions are inequivalent in general. In the present case, however, the large values of \\Ngal\\ that we consider prevent us from distinguishing between these two predictions, each of which provides an acceptable description of the data.  Additionally, we performed other tests to check the consistency of our results. We showed that the statistics of the luminosity gap (the difference in absolute magnitudes of the BSG and BCG) in the Mr20 catalog are consistent with the predictions of order statistics (equation~\\ref{pgapfix}). In Appendix~B we argued that statistical tests for the mean BCG luminosity based on fixed $N$ are to be preferred over those based on, say, fixed total group luminosity $L_{\\rm tot}$, showing how the latter may lead to biased conclusions. We also emphasized that, in order to perform unbiased tests, it is crucial to use accurate descriptions of the bright tail of the galaxy luminosity function.  However, we showed that, at the level of two-point statistics, the BCG and BSG luminosity distributions are inconsistent with their having being drawn from a universal luminosity function (i.e., one that is the same for all groups).  We argued that because the luminosity mark correlation function differs significantly from unity on large scales (Figure~\\ref{shuffledMarks}) the luminosity function of groups must depend on group properties (e.g. mass). This was not at all obvious from the (more traditional) one-point analysis of the group and BCG luminosity functions themselves.  In work in progress, we show that the extreme values assumption together with a universal luminosity distribution \\pgal\\ predicts a halo mass dependent BCG luminosity function $g_1(L|m) = \\sum_N p(N|m) g_1(L|N)$ if the distribution of \\Ngal\\ depends on mass $m$. This would have immediate consequences for implementations of the Halo Model such as the Halo Occupation Distribution (HOD, Zehavi \\etal\\ 2005), since such analyses would considerably simplify upon imposing the extreme value restriction on the BCG luminosities. However, we have already seen that the assumption of universality is inconsistent with the observed mark correlation signal on large scales. One might then wonder if the correct solution is to replace $p_{\\rm gal}(L)\\to p_{\\rm gal}(L|m)$. In this case, the BCG luminosity function averaged over \\Ngal\\ is mass dependent because both $p_{\\rm gal}$ (and hence $g_1(L|N,m)$) and $p(N|m)$ depend on $m$. In this case it is no longer obvious that the extreme values hypothesis should be tested at fixed \\Ngal\\ as we argued earlier. Formulating the problem cleanly in this case is work in progress.  Additionally, as discussed in the Introduction, this behaviour of the BCG luminosities, in particular, the fact that their luminosity functions are narrower than those of all galaxies, potentially allows BCGs to be used for consistency checks of standard candle constraints on the luminosity-distance relation. As a final remark, we note that, while BCGs have brighter and narrower luminosity distributions than randomly picked galaxies, the distributions of the BSG luminosities are even narrower than those of the BCGs. From the point of view of using them as standard candles, an interesting trade-off then arises: while it is the BCG which is most easily observed out to large distances, the BSG is a more standard candle. It will be interesting to see to what extent these ideas can be implemented in upcoming cluster surveys."
    },
    "1107/1107.2242_arXiv.txt": {
        "abstract": "{Open clusters are ideal test particles for studying the chemical evolution of the Galactic disc. However, the number and accuracy of existing high-resolution abundance determinations, not only of [Fe/H], but also of other key elements, remains largely insufficient.} {We attempt to increase the number of Galactic open clusters that have high quality abundance determinations, and to gather all the literature determinations published so far.} {Using high-resolution (R$\\sim$30000), high-quality (S/N$\\geq$60 per pixel), we obtained spectra for twelve stars in four open clusters with the fibre spectrograph FOCES, at the 2.2 Calar Alto Telescope in Spain. We employ a classical equivalent-width analysis to obtain accurate abundances of sixteen elements: Al, Ba, Ca, Co, Cr, Fe, La, Mg, Na, Nd, Ni, Sc, Si, Ti, V, and Y. We derived oxygen abundances derived by means of spectral synthesis of the 6300~\\AA\\  forbidden line.} {We provide the first determination of abundance ratios other than Fe for NGC~752 giants, and ratios in agreement with the literature for the Hyades, Praesepe, and Be~32. We use a compilation of literature data to study Galactic trends of [Fe/H] and [$\\alpha$/Fe] with Galactocentric radius, age, and height above the Galactic plane. We find no significant trends, but some indication for a flattening of [Fe/H] at large $R_{gc}$, and for younger ages in the inner disc. We also detect a possible decrease in [Fe/H] with $|$z$|$ in the outer disc, and a weak increase in [$\\alpha$/Fe] with $R_{gc}$.} {} ",
        "introduction": "\\label{sec1}  Open clusters (hereafter OC) are ideal {\\em test particles} for studying the evolution of metallicity with time, inferring the so-called {\\em age-metallicity relation}, and with Galactocentric radius, the {\\em metallicity gradient}, measuring in the Galactic disc. Their properties can be determined with smaller uncertainties than for field stars, since they are coeval group of stars at the same distance that have a homogeneous chemical composition. Unfortunately, of the $\\simeq$1700 known OC \\citep[e.g.][]{dias2002}, only $\\simeq$140 possess some metallicity determination, mostly obtained from photometric indicators, such as Washington or Str\\\"omgren photometry \\citep[see][and references therein]{twarog1997,chen2003} and low-resolution spectroscopy \\citep[e.g.][]{friel1993,friel02}.  The most accurate way to determine the chemical abundances is to analyse high-resolution spectroscopy. It allows us to investigate not only metallicity, but also abundance ratios -- with respect to iron or hydrogen -- of other chemical species such as $\\alpha$-elements, $s$-process elements, and $r$-process elements, which are synthesised in different environments and on different timescales (e.g. SNe Ia, SNe II, giants, supergiants, etc). In the past few years, a number of research groups have addressed the challenge of increasing the number of OC with chemical abundances determined from high-resolution spectroscopy \\citep[e.g.][]{sestito2004,dorazi2006,sestito2006,bragaglia2008,pace2008,dorazi2009,friel2010,pace2010,pancino2010a,jacobson2011}. However, the number of OC with chemical abundances determined with this technique is still small (see Section \\ref{sec6}), and significant uncertainties remain in the determinations of both the metallicity gradient and the age-metallicity relation, which are the fundamental ingredients of chemical evolution models.  \\begin{table*} \\begin{minipage}[t]{17.5cm} \\caption{Observing logs and programme star properties.} \\label{obslog} \\centering \\renewcommand{\\footnoterule}{} \\begin{tabular}{l c c c c c c c c c c c c} \\hline\\hline Cluster & Star & $\\alpha_{2000}$ & $\\delta_{2000}$ &  B & V & R  & I\\footnote{All I magnitudes are in the Johnson system (I$_J$) with the exception of those of the Be~32 stars which are in the Cousins system (I$_C$).} & K$_{S}$ & n$_{exp}$ & t$^{tot}_{exp}$ & S/N$^{tot}$\\\\ &  & (hrs) & (deg) & (mag) & (mag) &(mag) & (mag)& (mag)&  & (sec)& \\\\ \\hline Be 32\\footnote{Star names from \\citet{Richtler2001}; Coordinates, B, V \\& I$_C$ magnitudes from \\citet{dorazi2006}; K$_S$ magnitudes from 2MASS.} & 0456 & 06:58:08.2 & +06:24:19.6 & 14.76 & 13.67 & --- & 12.53 & 11.03 & 7 & 18900 & 60 \\\\ & 1948 & 06:58:04.2 & +06:27:17.1 & 14.50 & 13.37 & --- & 12.20 & 10.68 & 6 & 16200 & 70 \\\\ NGC 752\\footnote{Star names from \\citet{Heinemann1926}; Coordinates from \\citet{hog2000}; B \\& V magnitudes from \\citet{jennens75}; K$_S$ magnitudes from 2MASS.} & 001 & 01:55:12.6 & +37:50:14.6 & 10.47 & 9.51 & --- & --- & 7.23 & 4 & 2400 & 160 \\\\ & 208 & 01:57:37.6 & +37:39:38.1 & 10.04 & 8.97 & ---  & --- & 6.41 & 4 & 2400 & 180 \\\\ & 213 & 01:57:38.9 & +37:46:12.5 & 10.08 & 9.07 & --- & --- & 6.68 & 3 & 1800 & 80 \\\\ & 311 &01:58:52.9 & +37:48:57.3 & 10.11 & 9.07 & --- & --- & 6.64 & 4 & 2400 & 100 \\\\ Hyades\\footnote{Star names from \\citet{vanBueren1952}; Coordinates from \\citet{perryman1997}; B, V, R \\& I$_J$ magnitudes from \\citet{Johnson1966}; K$_S$ magnitudes from 2MASS.} & 028 ($\\gamma$ tau) & 04:19:47.6 & +15:37:39.5 & 4.64 & 3.65 & 2.92 & 2.45 & 1.52 & 2 & 120 & 560 \\\\ (Mel~25) & 041 ($\\delta$ tau) & 04:22:56.1 & +17:32:33.0 & 4.75 & 3.76 & 3.03 & 2.56 & 1.64  & 3 & 180 & 450 \\\\ & 070 ($\\epsilon$ tau) & 04:28:37.0 & +19:10:49.5 & 4.55 & 3.54 & 2.81& 2.31 & 1.42  & 3 & 180 & 270 \\\\ Praesepe \\footnote{Star names from \\citet{kleinWassink1927}; Coordinates from \\citet{perryman1997}; B, V \\& R magnitudes from \\citet{coleman82}; I$_J$ magnitudes from \\citet{mendoza1967,Johnson1966}, K$_S$ magnitudes from 2MASS.} & 212 & 08:39:50.7 & +19:32:27.0 & 7.53 & 6.58 & 5.87 & 5.38 & 4.39 & 4 & 240 & 165 \\\\ (NGC~2632) & 253 & 08:40:06.4 & +20:00:28.1 & 7.35 & 6.38 &5.67 & 5.20 & 4.20 & 4 & 240 & 215 \\\\ (M~44)   & 283 & 08:40:22.1 & +19:40:11.9 & 7.42 & 6.41 & 5.68 & 5.21 & 4.18 & 2 & 120 & 150 \\\\ \\hline\\hline \\end{tabular} \\end{minipage} \\end{table*}  In this paper, the second of a series initiated by \\citet[hereafter Paper I]{pancino2010a}, we present high quality and homogeneous measurements of chemical abundances for red clump stars in four OC: Be~32, NGC~752, Hyades, and Praesepe. The Hyades is the nearest OC and its four known red giants have been widely studied \\citep{schuler2009,mishenina2007,fulbright2007,schuler2006,mishenina2006,boyarchuk2000,luck1995}, hence it provides a very good reference frame to compare our abundances with the literature. Both NGC~752 and Praesepe have been well-studied, but all information about their chemical composition is based mainly on their main-sequence stars \\citep[e.g.][]{pace2008,an2007,sestito2004,burkhart1998,hobbs1992}. To our knowledge, there have been no recent measurements of the chemical abundance of their giants from high-resolution spectroscopy. Finally, Be~32 has been the subject of some studies \\citep[e.g.][]{Richtler2001,friel2010,bragaglia2008,dorazi2006}. The properties and previous studies of each cluster is described in more depth in Section \\ref{sec5}.  This paper is structured as follows: observations and data reduction are described in Section \\ref{sec2}; equivalent-width measurements are presented in Section \\ref{sec3}, together with the abundance analysis and its uncertainties; results are compared with the literature in Sections \\ref{sec5}, \\ref{sec6}, and \\ref{sec7}; and finally our main conclusions are summarised in Section \\ref{sec8}. ",
        "conclusions": "\\label{sec8}  We have enlarged our sample of homogeneous high-resolution abundance measurements from the five clusters of Paper~I to a total of nine, analysing here spectra of red clump giants in the Hyades, Praesepe, NGC~752, and Be~32. Our main results can be summarized as follows:  \\begin{itemize} \\item{We provide the first high-resolution based abundance ratios \\citep[other than {[}Fe/H{]}, see][]{sestito2008} for NGC~752, which turned out to be mostly of solar composition;} \\item{We have presented the abundance ratios of Praesepe red clump giants, which appear to solve a puzzling dichotomy of literature determinations for stars of different evolutionary stages;} \\item{We have found that our abundance ratios for the Hyades and Berkeley~32 are in good agreement with other literature determinations;} \\item{We have confirmed the absence of light elements (anti-)correlations in the OC studied so far.} \\end{itemize}  We have updated our compilation of previous literature data for 57 clusters of Paper~I to a total of 89 clusters presented here. With this updated compilation and our homogeneous measurements in hand, we have investigated Galactic trends in [Fe/H] (and [$\\alpha$/Fe]) with age, Galactocentric radius, and height above the Galactic plane. Our findings are in substantial agreement with other similar investigations, where the abundance gradient appears to indeed flatten out outside $R_{gc}\\simeq$12.5~kpc, and the inner disc slope appears to flatten for younger ages as well, although the age bins are not too well-sampled. At the same time, [$\\alpha$/Fe] shows a weak increase with $R_{gc}$. No significant gradients are observed with $|$z$|$ or age, except for a weak tendency of [Fe/H] to decrease with increasing $|$z$|$ and decrease with age in the outermost disc annulus studied. None of our measured weak trends have any significance above 1\\,$\\sigma$. Larger samples of homogeneous data are still necessary to investigate the existence of any dependence on age and $|$z$|$ in the Galactic disc."
    },
    "1107/1107.3843_arXiv.txt": {
        "abstract": "{We present multi-configuration Breit-Pauli distorted-wave photoionization (PI) cross sections and radiative recombination (RR) and dielectronic recombination (DR) rate coefficients for the first six krypton ions.  These were calculated with the AUTOSTRUCTURE code, using semi-relativistic radial wavefunctions in intermediate coupling.  Kr has been detected in several planetary nebulae (PNe) and H~II regions, and is a useful tracer of neutron-capture nucleosynthesis.  PI, RR, and DR data are required to accurately correct for unobserved Kr ions in photoionized nebulae, and hence to determine elemental Kr abundances.  PI cross sections have been determined for ground configuration states of Kr$^0$--Kr$^{5+}$ up to 100~Rydbergs.  Our Kr$^{+}$ PI calculations were significantly improved through comparison with experimental measurements.  RR and DR rate coefficients were respectively determined from the direct and resonant PI cross sections at temperatures (10$^1-10^7$)$z^2$~K, where $z$ is the charge.  We account for $\\Delta n=0$ DR core excitations, and find that DR is the dominant recombination mechanism for all but Kr$^+$ at photoionized plasma temperatures.  Internal uncertainties are estimated by comparing results computed with three different configuration-interaction expansions for each ion, and by testing the sensitivity to variations in the orbital radial scaling parameters.  The PI cross sections are generally uncertain by 30--50\\% near the ground state thresholds.  Near 10$^4$~K, the RR rate coefficients are typically uncertain by less than 10\\%, while those of DR exhibit uncertainties of factors of 2 to 3, due to the unknown energies of near-threshold autoionizing resonances.  With the charge transfer rate coefficients presented in the third paper of this series, these data enable robust Kr abundance determinations in photoionized nebulae for the first time, providing a new tool for studying heavy element enrichments in PNe and for investigating the chemical evolution of trans-iron elements.} ",
        "introduction": "\\label{intro}  Recent measurements of neutron(\\emph{n})-capture element (atomic number $Z>30$) emission lines in ionized astrophysical nebulae have spurred laboratory astrophysics investigations of these species.  Trans-iron elements were first detected in planetary nebulae (PNe) 35 years ago \\citep{treffers76}, but it was not until nearly two decades later that their emission lines were identified \\citep{pb94}.  The detection of these elements marked a new era in the field of PNe, since the progenitor stars of these objects can produce trans-iron elements via slow \\emph{n}-capture nucleosynthesis (the \\emph{s}-process) during the asymptotic giant branch (AGB) phase of evolution.  Abundance determinations of PNe complement spectroscopic investigations of AGB stars, enabling previously unexplored details of \\emph{s}-process nucleosynthesis to be revealed \\citep[e.g., see][]{sterling08, sterling11b}.  Indeed, many of the \\emph{n}-capture elements that are accessible in nebular spectra, such as Se, Kr, and Xe, cannot be detected in cool giants such as AGB stars (and in fact had not previously been studied in any of their sites of origin).  The pioneering work of \\citet{pb94} inspired other groups \\citep{dinerstein01a, dinerstein01b, sharpee07} to identify previously unrecognized nebular emission lines as transitions of trans-iron elements.  In spite of the low cosmic abundances of \\emph{n}-capture elements \\citep{asplund09}, deep ultraviolet \\citep{sterling02, sterling03, williams08}, optical \\citep[e.g.,][]{hyung01, sharpee03, liu04, zhang05, sharpee07, sterling09, otsuka10b, otsuka11} and near-infrared spectroscopy \\citep{likkel06, sterling07, sterling08} of several Galactic and Local Group PNe revealed spectral features of these species.  Abundances of these elements have now been estimated in approximately 100 PNe, allowing for an analysis of \\emph{s}-process elemental enrichments in a statistically meaningful sample of objects.  These detections are not limited to PNe.  Indeed, \\emph{n}-capture element emission lines have been detected in a variety of other environments, including H~II regions \\citep{aspin94, lumsden96,luhman98, baldwin00, puxley00, okumura01, blum08, roman-lopes09} and starburst galaxies \\citep{vanzi08}.  These detections demonstrate the promise of nebular spectroscopy as a tool to study \\emph{n}-capture nucleosynthesis in low-mass stars, the chemical evolution of trans-iron elements in the Universe, and the heavy-element nucleosynthetic histories of other galaxies.  Kr is one of the most readily detected \\emph{n}-capture elements in ionized nebulae, along with Se and Xe, with relatively strong transitions in both the optical \\citep{pb94} and near-infrared \\citep{dinerstein01a}.  It has been detected in dozens of Galactic PNe and H~II regions, as well as a handful of extragalactic nebulae \\citep[e.g.,][]{wood07, otsuka11}.  Beyond the relative ease of detection, Kr is a sensitive probe of \\emph{s}-process nucleosynthesis.  It can be highly enriched by the \\emph{s}-process, due in part to the $^{86}$Kr isotope, which has a magic number of neutrons and acts as a bottleneck in the \\emph{s}-process path.  Isotopes with magic numbers of neutrons have very small \\emph{n}-capture cross sections and cause the peaks seen in the element-by-element \\emph{s}-process enrichment distribution near $Z=40$, 56, and 82 \\citep[e.g.,][]{busso99, burris00, sneden08}.  Typically, \\emph{s}-process nucleosynthesis enriches Kr by a significantly larger factor than Se \\citep{sterling08, karakas09}, which does not have a stable isotope with a magic number of neutrons.  In addition, because it is a noble gas, Kr is not depleted into dust grains, and hence its gas-phase abundance is representative of the total elemental abundance.  However, nebular spectroscopy requires a foundation of atomic data to determine both ionic abundances and ionization balance solutions to correct for unobserved ions.  Transition probabilities and effective collision strengths have been determined for many detected \\emph{n}-capture element ions \\citep{biemont86a, biemont86b, biemont87, biemont88, biemont95, schoning97, sb98}, enabling their ionic abundances to be derived.  However, one of the primary hurdles to nebular studies of trans-iron elements is that the ionization equilibria of these species cannot be accurately solved due to the lack of photoionization and recombination data for nearly all \\emph{n}-capture element ions.  Motivated by this clear need for atomic data, we have embarked on a program using both theoretical and experimental methods to determine photoionization (PI) cross sections and rate coefficients for radiative recombination (RR), dielectronic recombination (DR), and charge transfer (CT).  Collisional ionization and high-temperature DR are of negligible importance for ionization balance solutions in photoionized nebulae such as PNe and H~II regions.  This paper is the second in a series presenting atomic data calculations for \\emph{n}-capture element ions.  In the first \\citep{sterling11b}, multi-configuration Breit-Pauli (MCBP) distorted-wave PI cross sections and RR and DR rate coefficients were given for the first six ions of Se.  We furnish CT rate coefficients for the first five ions of Ge, Se, Br, Kr, Rb, and Xe in the third paper \\citep{sterling11c}, computed in the Demkov and Landau-Zener approximations.  In this study, we present MCBP distorted-wave PI cross sections and RR and DR data for the first six Kr ions, calculated with the atomic structure code AUTOSTRUCTURE \\citep{badnell86, badnell97, badnell11}.  We do not investigate more highly-charged Kr species, since PN central stars and massive young stars (the ionizing sources of H~II regions) are not sufficiently hot to significantly photoionize Kr more than six times \\citep[Kr$^{6+}$ has an ionization potential of 111~eV, and PN central stars generally have effective temperatures less than 200\\,000~K;][]{napiwot99, phillips03, villaver07}.  Photoionization and recombination studies of low-charge Kr ions are few and far between, particularly in terms of theoretical calculations.  On the other hand, neutral Kr has been the subject of numerous experimental \\citep[e.g.,][and references therein]{henke67, lang75, marr76, west76, samson89, samson91, sorokin00} and theoretical \\citep[e.g.,][]{amusia71, miller77, parpia84, tulkki92} PI studies.  However, neutral Kr is a trace species in PNe and H~II regions, and therefore its PI properties are not expected to play a significant role in Kr ionization balance solutions.  In terms of recombination, little is known about the ions of interest in our work.  \\citet{mattioli06} compiled RR and DR rate coefficients for all Kr ions, but the data for low-charge ions are rough estimates.  For these species, DR rate coefficients were computed using the Burgess formula \\citep{burgess65}, as modified by  \\citet{merts76}, which is appropriate for collisionally-ionized plasmas (indeed, the work of Mattioli et al.\\ was inspired by the utility of Kr as a coolant for magnetically-confined fusion plasmas), but is known to badly underestimate the rate coefficient at the low temperatures of photoionized plasmas.  Their RR rate coefficients for low-charge Kr ions are approximate, with RR into the valence and (hydrogenic) excited shells treated separately.  Experimental measurements of absolute PI cross sections provide the most useful comparisons to our calculations.  Concurrent to our theoretical investigations, we have experimentally measured PI cross sections for Se and Xe ions \\citep{esteves09, esteves10, sterling11a, esteves11a, esteves11b} at the Advanced Light Source synchrotron radiation facility, for the purpose of constraining our calculations.  PI investigations were previously carried out for Kr$^{3+}$ to Kr$^{5+}$ \\citep{lu06a, lu06b, lu_thesis}, and more recently for Kr$^+$ \\citep{bizau11}.  We compare our computed PI cross sections to these measurements to benchmark and gauge the accuracy of our direct PI cross sections.  Because no comprehensive studies of the PI and recombination properties of low-charge Kr ions in the photoionized regime existed prior to our study, we provide estimates for the internal uncertainties of our computed atomic data.  These uncertainties can affect ionization equilibrium calculations, and hence elemental abundance determinations.  Beyond comparisons to experimental PI cross-section measurements, we computed data using three different configuration-interaction (CI) expansions for each Kr ion to gauge the sensitivity of the results to CI effects.  Moreover, we tested the effects of other parameters and assumptions in our calculations (e.g., whether the radial wavefunctions were orthogonalized, the radial scaling parameters for continuum orbitals, etc.) on the resulting data.  In the remainder of this paper, we present the computed PI cross sections and RR and DR rate coefficients for low-charge Kr ions.  In Sect.~\\ref{calcs}, we briefly describe the calculations and methodology for each of these processes, and detail the calculated electronic structure for each ion.  The resulting data are presented in Sect.~\\ref{results}, along with estimated uncertainties and comparisons to previous studies.  We summarize our investigation in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We have presented MCBP calculations of distorted-wave photoionization cross sections over the photon energy range 0--100~Ryd, and radiative and dielectronic recombination rate coefficients at temperatures (10$^1-10^7$)$z^2$~K for the first six Kr ions.  The results of all of our calculations are available at the CDS, as well as at http://www.pa.msu.edu/astro/atomicdata/kr\\_data.tar.gz.  The RR and DR rate coefficients have been fit to analytical functions, with fit coefficients given in Tables~\\ref{rrlowtfits}, \\ref{rrhitfits}, and \\ref{drfits}.  Kr has been detected in the spectra of planetary nebulae and H~II regions, and its abundance is a sensitive tracer of enrichments by \\emph{n}-capture nucleosynthesis.  But in order to derive elemental Kr abundances from those of the observed ions, unseen ionization stages must be accounted for via ionization equilibrium solutions.  Our atomic data determinations, along with charge transfer rate coefficients presented in the third paper of this series \\citep{sterling11c}, enable much more accurate Kr ionization balance computations than previously possible.  This study represents the first comprehensive investigation of the photoionization and recombination properties of low-charge Kr ions specifically focused on the photoionized plasma regime.  As such, we have taken significant efforts to provide realistic internal uncertainties for our results, both through comparison to experimental PI measurements, and by testing the sensitivity of our data to changes in the CI expansion and orbital radial scaling parameter values.  We find typical uncertainties of 30--50\\% for our computed PI cross sections, though comparison to experimental measurements indicates that the Kr$^0$ and Kr$^{3+}$ cross sections have larger uncertainties of roughly a factor of two near their ionization thresholds.  We note that our Kr$^+$ PI calculations were significantly improved after comparison with experimental measurements, from accuracies of a factor of 2.5 near the ground state threshold to within $\\sim$40\\%.  This demonstrates the utility and importance of experimental absolute PI cross-section measurements for constraining and guiding theoretical calculations.  Our RR rate coefficients are more precisely determined, with uncertainties less than 10\\% except for the Kr$^+$ target (50--60\\%).  On the other hand, the unknown energies of low-lying autoionizing states lead to significantly larger uncertainties in DR rate coefficients, typically factors of 2--3.  These uncertainties are magnified by the fact that DR is the dominant form of recombination for low-charge Kr ions at the temperatures of photoionized astrophysical nebulae, with the rate coefficients exceeding those of RR by as much as two orders of magnitude.  The dominance of DR over RR in photoionized nebulae is common for heavy element ions \\citep[e.g.,][]{badnell06a, altun07, nikolic10, sterling11b}.  The estimated uncertainties in these atomic data will play an important role in our forthcoming study of the Se and Kr ionization balance in planetary nebulae.  We will numerically simulate the Se and Kr ionization balance for a wide range of nebular parameters to derive new analytical corrections for unobserved ions.  Using our estimated uncertainties in the PI cross sections and RR and DR rate coefficients, we will run Monte Carlo simulations to test how sensitive elemental abundance determinations are to these uncertainties, thereby illuminating the atomic processes and ionic systems that require additional theoretical and/or experimental investigation."
    },
    "1107/1107.3558_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  We are still a long way from observationally mapping active or future cluster-forming regions on a one-by-one basis in distant galaxies.  Even at ALMA resolution (0.1 arcsec), to map clumps of cold dense molecular gas on a pc-scale is attainable only in galaxies less distant than $\\simeq 2$\\,Mpc.  In contrast, properties of {\\it systems} of star clusters, such as the distribution functions of cluster ages, masses and radii, are retrievable out to distances of at least 20\\,Mpc.  These distribution functions are shaped by both the dynamical evolution and the formation conditions of star clusters.  Therefore, evolving model cluster systems for a wide range of initial conditions, and comparing their predicted distribution functions to their  observed counterparts, constitutes a powerful tool to probe into the physics of cluster formation.  That `macroscopic' approach has the tremendous advantage to alleviate the need for a spatial resolution which is yet to be achieved in galaxies more distant than a few Mpc (`microscopic' approach).  Systems of young clusters are -- obviously -- those best suited to uncover cluster-forming region properties.  At ages less than 100\\,Myr, a key-driver of cluster evolution is violent relaxation -- namely the dynamical response of embedded clusters to the expulsion of their residual star-forming gas due to massive star activity \\citep[e.g.][]{2001MNRAS.323..988G}.  During violent relaxation, clusters expand and lose stars (infant weight-loss) as a result of the binding gaseous mass loss.  Initial conditions of this phase leave therefore signatures in the cluster age, mass and radius distribution functions \\citep[e.g.][]{2008ApJ...678..347P, 2009ApJ...690.1112P}.  In this contribution, I highlight how the time-invariant shape of the {\\it young} cluster mass function sheds light on the mass-radius relation of cluster-forming regions (hereafter CFRg).  I show that this relation is one of constant mean volume density, that is, $r_{CFRg} \\propto m_{CFRg}^{1/3}$ with $r_{CFRg}$ and $m_{CFRg}$ the CFRg radius and mass, respectively.  Furthering this argument, I demonstrate: (i) why the young cluster mass function is steeper than the molecular cloud mass function, and (ii) why there is a massive star formation limit in the mass-size space of molecular structures. ",
        "conclusions": "}  The analysis of the impact exerted by a tidal field upon star clusters which have expelled their residual star-forming gas has been carried out.  It shows that the time-invariant shape of the young star cluster mass function is robustly reproduced if the mean {\\it volume} density of cluster-forming regions (CFRgs) is constant.  If the mass-radius relation of CFRgs were one of constant mean {\\it surface} density, it would lead to the preferential removal of high mass-clusters, which contradicts observational results.  These trends follow from how CFRg volume density and CFRg mass scale with each other.  This stellar-dynamics-based  finding is independently confirmed by studies mapping star formation activity in molecular clumps, molecular clouds and galaxies, which all show that star formation is driven by a volume (number) density threshold.  That star formation takes place in gas denser than a given number density threshold, $n_{{\\rm H_2},th}$, and that molecular clumps are characterized by density gradients allow us to define two distinct regions in molecular clumps: a central dense CFRg, and an outer envelope inert in terms of star formation due to $n_{{\\rm H_2}}<n_{{\\rm H_2},th}$.  As a result, properties (e.g. mass, radius, mean volume densities) of CFRgs and molecular clumps get dissociated.  Building on that picture, I put forward an original explanation for why the mass function of young star clusters is steeper than that of molecular clumps and clouds.  The mass-radius relation of molecular clouds and, possibly, of star-forming clumps, is one of constant mean surface density.  This contrasts with the constant mean volume density inferred for CFRgs.  This difference in slope in their respective mass-radius relations steepens the mass function of CFRgs compared to that of their host-clumps because clumps of higher mass have a lower mean volume density (hence a smaller fraction of star-forming gas) at constant mean surface density.  Finally, the same model is successfully applied to understand the origin of the massive star formation limit (MSF) in the mass-size space of molecular structures. This is shown to be consistent with a {\\it threshold in star-forming gas mass} beyond which the star-forming gas reservoir is large enough to allow the formation of massive stars.  Specifically, the MSF limit is consistent with the formation of a $10\\,{\\rm M}_{\\odot}$-star out of its individual density peak with $n_{{\\rm H_2}, th} \\simeq 10^5\\,cm^{-3}$, or with the formation of a $10\\,{\\rm M}_{\\odot}$-star as a CFRg member with $n_{{\\rm H_2}, th} \\simeq 10^4\\,cm^{-3}$. \\\\    Slides of oral presentations related to the topics discussed in this contribution are available at \\href{http://www.astro.uni-bonn.de/~gparm/talks.html}{http://www.astro.uni-bonn.de/$\\sim$gparm/talks.html}   \\small  %"
    },
    "1107/1107.1241_arXiv.txt": {
        "abstract": "We investigate the general properties of expanding cosmological models which generate scale-invariant curvature perturbations in the presence of a variable speed of sound. We show that in an expanding universe, generation of a super-Hubble, nearly scale-invariant spectrum of perturbations over a range of wavelengths consistent with observation requires at least one of three conditions: (1) accelerating expansion, (2) a speed of sound faster than the speed of light, or (3) super-Planckian energy density. ",
        "introduction": "Observations of both the Cosmic Microwave Background (CMB) and large-scale structure (LSS) are compatible with a spectrum of nearly scale-invariant and Gaussian primordial perturbations with correlations on super-Hubble length scales \\cite{Komatsu:2010fb}, in agreement with the predictions of the simplest inflationary models. However, inflation is not the only way to generate cosmological perturbations. Alternatives that have been put forward range from contracting scenarios where the universe needs to go through a singular or non-singular bounce (see \\cite{Khoury:2011ii} and references there), string gas cosmology, in which the universe is initially static \\cite{Nayeri:2005ck}, varying fundamental speed of light \\cite{Albrecht:1998ir}, or a rapidly varying speed of sound \\cite{ArmendarizPicon:2003ht,ArmendarizPicon:2006if,Magueijo:2008pm, Bessada:2009ns, Magueijo:2009zp}. All of the alternatives so far suggest that to produce the observed primordial power spectrum, any route but inflation requires an understanding of gravity beyond general relativity. This is either due to violation of the Null Energy Condition, a singular bounce, or the breaking of Lorentz invariance due to a preferred frame of reference to explain  varying speed of light or super-luminal sound speed \\cite{Babichev:2007dw}. (It is worth noting that inflation itself, while more successful than its alternatives on many fronts, still has its own challenges to overcome, ranging from setting the correct initial conditions \\cite{Vachaspati:1998dy} to relying on the existence of a scalar field which has not yet been confirmed by any fundamental theory.) It is interesting to ask: what is the most general conclusion that can be drawn from the observed spectrum of primordial perturbations in the universe?   In this paper, we study general properties of the production of a scale-invariant two point function. We use the simple framework  put forward in Ref. \\cite{Khoury:2008wj} and later used in Ref. \\cite{Khoury:2010gw} that makes it simple to study scenarios  that have different background evolution but result in the same second-order gravitational action for quantum perturbations. The format of the paper is the following: In Section \\ref{setup} we review the framework developed in Ref. \\cite{Khoury:2008wj}. In Section \\ref{nogo} we prove that to produce enough scale invariant modes  on super-Hubble scales in an expanding universe one of the following conditions must be met: (1) accelerating expansion, (2) superluminal sound speed, or (3) super-Planckian energy density. Section \\ref{conclusion} contains discussion and conclusions. ",
        "conclusions": "\\label{conclusion}  The universe is observed to have a spectrum of nearly scale-invariant density perturbations over about three decades in wavelength, which were at scales larger than the Hubble length  at early times. In this paper, we have shown that generating cosmological perturbations consistent with observation in an expanding universe requires at least one of: (1) accelerated expansion, (2) superluminal sound speed, or (3) super-Planckian energy density. We note that this applies to the curvaton mechanism \\cite{Lyth:2001nq} as well, since the freezeout horizon for a free scalar ``spectator'' field shrinks as $\\tau$ ($a \\propto 1/\\tau)$ to produce scale-invariant perturbations. Therefore super-Hubble curvaton fluctuations directly require inflation. It is important to note that scale invariance alone does not require inflation or ``tachyacoustic'' \\cite{Bessada:2009ns} evolution: a key point is that even with sub-luminal sound speed, super-Hubble perturbations can be generated with a growing Hubble horizon, as long as the freezeout horizon is shrinking and $R_\\zeta > R_H$ \\cite{Khoury:2010gw}. However, one cannot generate three decades of modes in this fashion without super-Planckian energy densities. Perturbations could also be generated on sub-Hubble scales by a period of non-inflationary expansion \\cite{Khoury:2010gw,Khoury:2011ii}, and only later redshifted to super-Hubble scales by a subsequent period of inflation. This would be consistent with our bound. Recent work has, however, shown why seeding  scale-invariant fluctuations with sub-luminal sound speed is challenged by the breakdown of weak coupling and therefore perturbation theory \\cite{Baumann:2011dt}. This issue is quite generic and can arise even in non-expanding scenarios (see \\cite{Khoury:2011ii}). Similarly, string gas cosmology requires accelerating expansion to exit from the initial, static phase, but perturbations are not generated by the usual inflationary mechanism.  Our bound also does not apply to contracting cosmologies. However, in the case of an expanding cosmology, we have presented a simple and general argument showing why either a superluminal sound speed or a period of inflation is required for the successful generation of a scale-invariant, super-Hubble spectrum of perturbations."
    },
    "1107/1107.2901.txt": {
        "abstract": "This paper details the solar neutrino analysis of the 385.17-day Phase-III data set acquired by the Sudbury Neutrino Observatory (SNO).   An array of $^3$He proportional counters was installed in the heavy-water target to measure precisely the rate of neutrino-deuteron neutral-current interactions.  This technique to determine the total active $^8$B solar neutrino flux was largely independent of the methods employed in previous phases.  The total flux of active neutrinos was measured to be $\\snoncfluxunc$ $\\times~10^{6}$~cm$^{-2}$~s$^{-1}$, consistent with previous measurements and standard solar models.  A global analysis of solar and reactor neutrino mixing parameters yielded the best-fit values of  $\\Delta m^2 = \\snodmsquared$ and $\\theta = \\snothetaonetwo$. ",
        "introduction": "}  The Sudbury Neutrino Observatory (SNO) experiment~\\cite{NIM} has firmly established that electron-type neutrinos ($\\nu_e$) produced in the solar core transform into other active flavors while in transit to the Earth~\\cite{snocc,snonc,snodn,longd2o,nsp,ncdprl}.  This direct observation of neutrino flavor transformation, through the simultaneous observation of the disappearance of $\\nu_e$ and the appearance of other active neutrino types, confirmed the total solar neutrino flux predicted by solar models~\\cite{BP01,TC}, and explained the deficit of solar neutrinos that was seen by other pioneering experiments~\\cite{cl,sage,gallex,SK,gno}.  The SNO results, when combined with other solar neutrino experiments and reactor antineutrino results from the KamLAND experiment~\\cite{kam}, demonstrated that neutrino oscillations~\\cite{pontecorvo,wolfenstein,ms} are the cause of this flavor change.  In the first two phases of the SNO experiment the determination of the total active $^8$B solar neutrino flux and its $\\nu_e$ component required a statistical separation of the Cherenkov signals observed by the detector's photomultiplier tube (PMT) array.  In the third phase of the experiment, an array of $^3$He proportional counters~\\cite{ncdcountersnim} was deployed in the detector's heavy-water target.  The neutron signal in the inclusive total active neutrino flux measurement was detected predominantly by this ``Neutral-Current Detection'' (NCD) array, and was separate from the Cherenkov light signals observed by the PMT array in the $\\nu_e$ flux measurement.  This technique to measure the total active  $^8$B solar neutrino flux was largely independent of the methods employed by SNO in previous phases.  The results from the third phase of the SNO experiment were reported in a letter~\\cite{ncdprl}, and confirmed those from previous phases.  We present in this article the details of this measurement and an analysis of the neutrino oscillation parameters.  In Sec.~\\ref{sec:overview} we present an overview of the SNO experiment and the solar neutrino measurement with the NCD array.  The Phase-III data set that was used in this measurement is described in Sec.~\\ref{sec:dataset}.  Details of the optical response of the PMT array and the reconstruction of its data are provided in Sec.~\\ref{sec:pmtresponse}.  The electronic and energy response of the NCD array will be discussed in Sec.~\\ref{sec:ncdresponse}.  The determination of the neutron detection efficiencies for both the PMT and the NCD arrays, which are crucial to the measurement of the total active solar neutrino flux, is presented in Sec.~\\ref{sec:neutron}.  The evaluation of backgrounds in the measurement is summarized in Sec.~\\ref{sec:backgrounds}.  Alpha decays in the construction materials of the NCD array were a non-negligible background for the detection of signal neutrons.  We developed an extensive pulse-shape simulation and applied it to understand the response of the NCD counters to these alpha decays.   The pulse-shape simulation model is presented in Sec.~\\ref{sec:pulsesim}.  In Sec.~\\ref{sec:sigex}, we discuss the analysis that determined the total active solar neutrino flux and the electron-type neutrino flux.   These measured fluxes, along with results from previous SNO measurements and other solar and reactor neutrino experiments, were then used in the determination of the neutrino mixing parameters as described in Sec.~\\ref{sec:physint}.   A description of the cuts we used to remove instrumental backgrounds in the NCD array data can be found in Appendix~\\ref{sec:apdxa}.  A discussion of the parameterization of nuisance parameters in the neutrino flux analysis is provided in Appendix~\\ref{sec:apdxb}. ",
        "conclusions": "} We have presented a detailed description of the SNO Phase-III results that were published in Ref.~\\cite{ncdprl}.  Neutrons from the NC reaction were detected predominantly by the NCD array.  The use of this technique, which was independent of the neutron detection methods in previous phases, resulted in reduced correlations between the CC, ES and NC fluxes and an improvement in the mixing angle uncertainty.  Several techniques to reliably calibrate the PMT and NCD arrays were developed and are detailed in this paper.  The presence of the NCD array changed the optical properties, and hence the energy response, of the PMT array.  Extensive studies and evaluation of the techniques used in calibrating the PMT array in previous phases were performed and reported.   Radioactive backgrounds associated with the NCD array, the \\dto\\ target and other detector components were precisely quantified by \\textit{in situ} and \\textit{ex situ} measurements.  These measurements provided the constraints for the respective nuisance parameters in the determination of the $\\nu_e$ and the total active neutrino fluxes.  The total flux of active neutrinos was measured to be $\\snoncfluxunc \\times 10^{6}$~cm$^{-2}$~s$^{-1}$, and was consistent with previous measurements and standard solar models.  A global analysis of neutrino mixing parameters using solar and reactor neutrino results yielded the best-fit values of the neutrino mixing parameters of $\\Delta m^2 = \\snodmsquared$ and $\\theta= \\snothetaonetwo$.  A detailed paper that describes an analysis of data combined from all three phases of SNO is in preparation.  %"
    },
    "1107/1107.1255_arXiv.txt": {
        "abstract": "The increasing size of cosmological simulations has led to the need for new visualization techniques.  We focus on Smoothed Particle Hydrodynamical (SPH) simulations run with the {\\small GADGET} code and describe methods for visually accessing the entire simulation at full resolution.  The simulation snapshots are rastered and processed on supercomputers into images that are ready to be accessed through a  web interface (GigaPan).  This allows any scientist with a web-browser to interactively explore simulation datasets in both in spatial and temporal dimensions, datasets which in their native format can be hundreds of terabytes in size or more. We present two examples, the first a static terapixel image of the MassiveBlack simulation, a {\\small P-GADGET} SPH simulation with 65 billion particles, and the second an interactively zoomable animation of a different simulation with more than one thousand frames, each a gigapixel in size. Both are available for public access through the GigaPan web interface.  We also make our imaging software publicly available. ",
        "introduction": "In the 40 years that $N$-body simulations have been used in Cosmology research, visualization has been the most indispensable tool.  Physical processes have often been identified first and studied  via images of simulations. A few examples are: formation of filamentary structures in the large-scale distribution of matter \\citep{JENKINS,COPLBERG,MILLENIUM-I,DOLAG}, growth of feedback bubbles around quasars \\citep{SIJACKI,FEEDBACK}; cold flows of gas forming galaxies  \\citep{COLDFLOW-DEKELA,KERES}, and the evolution of ionization fronts during the re-ionization epoch \\citep{BUBBLESHIN,BUBBLEZAHN}. The size of current and upcoming Peta-scale simulation datasets can make such visual exploration to discover new physics technically challenging. Here we present techniques that can  be used to display images at full resolution of datasets of hundreds of billions of particles in size.   Several implementations of visualization software for cosmological simulations already exist. IRFIT \\citep{IFRIT} is a general purpose visualization suite that can deal with mesh based scalar, vector and tensor data, as well as particle based datasets as points. YT \\citep{YT} is an analysis toolkit for mesh based simulations that also supports imaging. SPLASH \\citep{SPLASH} is a visualization suite specialized for simulations that use smoothed particle hydrodynamics(SPH) techniques. Aside from the CPU based approaches mentioned above, \\cite{GPUVIS} implemented a GPU based interactive visualization tool for SPH simulations.  The Millennium I \\& II simulations \\citep{MILLENIUM-I,MILLENIUM-II} have been used to test an interactive scalable rendering system developed by \\citet{FSW2009}. Both SPLASH and the Millennium visualizer support high quality visualization of SPH data sets, while IRFIT treats SPH data as discrete points.  Continuing improvements in  computing technology and algorithms are allowing  SPH cosmological simulations to be run with ever increasing numbers of particles.  Runs are now possible on scales which allow rare objects, such as quasars to form in a large simulation volume with uniform high resolution (see Section 2.1; and \\citet{DIMATTEO-PREP,DEGRAF-PREP,KHANDAI-PREP}).  Being able to scan through a vast volume and seek out the tiny regions of space where most of the activity is occurring, while still keeping the large-scale structure in context necessitates special visualization capabilities. These should be able to show the largest scale information but at the same be interactively zoomable.  However, as the size of the datasets quickly exceeds the capability of moderately sized in-house computer clusters, it becomes difficult to perform any interactive visualizations. For example, a single snapshot of the MassiveBlack simulation (Section 2.1) consists of 8192 files and is over 3 TB in size.  Even when a required large scale high resolution image has been rendered, actually exploring the data requires special tools.  The GigaPan collaboration \\footnote{http://www.gigapan.org} has essentially solved this problem in the context of viewing large images, with the GigaPan viewer enabling anyone connected to the Internet to zoom into and explore in real time images which would take hours to transfer in totality. The viewing technology has been primarily used to access large photographic panoramas, but is easily applicable to simulated datasets. A recent enhancement to deal with the time dimension, in the form of gigapixel frame interactive movies (GigaPan Time Machine\\footnote{http://timemachine.gigapan.org}) will turn out to give particularly novel and exciting results when applied to simulation visualization.    In this work we combine an off-line imaging technique together with GigaPan technology to implement an interactively accessible  visual probe of large cosmological simulations. While GigaPan is an independent project (uploading and access to the GigaPan website is publicly available), we release our  toolkit for the off-line visualization  as {\\small Gaepsi}% \\footnote{http://web.phys.cmu.edu/\\textasciitilde{}yfeng1/gaepsi% }, a software package aimed specifically at {\\small GADGET} \\citep{GADGET,GADGET2} SPH simulations.  The layout of our paper is as follows. In Section 2 we give a brief overview of the physical processes modeled in {\\small GADGET}, as well as describing two {\\small P-GADGET} simulations which we have visualized. In Section 3 we give details of the spatial domain remapping we employ to convert cubical simulation volumes into image slices. In Section 4, we describe the process of rasterizing an SPH density field, and in Section 5  the image rendering and layer compositing. In Section 6 we address the parallelism of our techniques and give measures of performance. In Section 7 we briefly describe the GigaPan and GigaPan Time Machine viewers and present examples screenshots from two visualizations (which are both accessible on the GigaPan websites). ",
        "conclusions": "We have presented a framework for generating and viewing large images and movies of the formation of structure in cosmological SPH simulations. This framework has been designed specifically to tackle the problems that occur with the largest datasets. In the generation of images, it includes parallel rasterization (for either shared and distributed memory) and adaptive pixel filling which leads to a well behaved filling rate at high resolution. For viewing images, the GigaPan viewers use hierarchical caching and cloud based storage to make even the largest of these datasets fully explorable at high resolution by anyone with an internet connection. We make our image making toolkit publicly available, and the GigaPan web resources are likewise publicly accessible."
    },
    "1107/1107.3585_arXiv.txt": {
        "abstract": "{Observations of pulsar timing provide strong constraints on scalar-tensor theories of gravity, but these constraints are traditionally quoted as limits on the microscopic parameters (like the Brans-Dicke coupling, for example) that govern the strength of scalar-matter couplings at the particle level in particular models. Here we present fits to timing data for several pulsars directly in terms of the phenomenological couplings (masses, scalar charges, moment of inertia sensitivities and so on) of the stars involved, rather than to the more microscopic parameters of a specific model. For instance, for the double pulsar PSR J0737-3039A/B we find at the 68\\% confidence level that the masses are bounded by $1.28 < m_{\\ssA}/m_{\\odot} < 1.34$ and $1.19 < m_{\\ssB}/m_{\\odot} < 1.25$, while the scalar-charge to mass ratios satisfy $|a_{\\ssA}| < 0.21$, $|a_{\\ssB}| < 0.21$ and $|a_{\\ssB}-a_{\\ssA}| < 0.002$. These constraints are independent of the details of the scalar tensor model involved, and of assumptions about the stellar equations of state. Our fits can be used to constrain a broad class of scalar tensor theories by computing the fit quantities as functions of the microscopic parameters in any particular model. For the Brans-Dicke and quasi-Brans-Dicke models, the constraints obtained in this manner are consistent with those quoted in the literature. } ",
        "introduction": "Although General Relativity (GR) has many applications in astrophysics and cosmology, incomplete knowledge about the gravitating bodies involved prevents using most of these as tests of the theory itself for the vast majority of systems outside of the solar system. Binary pulsars provide the rare exception to this rule, due to the great precision with which timing measurements allow the properties of their orbits to be inferred \\cite{BPrev,Damour:1991rd,PPK}. These orbital properties are parameterized by the inferred values of the Keplerian parameters (like semi-major axis, $\\frak{a}$, and eccentricity, $e$) that define the characteristics of the Newtonian orbit of the pulsar and its partner, together with its orientation in space. They are also parameterized by a suite of post-Keplerian (PK) parameters that describe observable slow, secular orbital changes over time, as well as relativistic time delays in the propagation of radio signals emitted by the pulsar.  \\subsection{The classic constraints}  GR predicts the values for these post-Keplerian parameters in terms of the underlying Keplerian parameters and the masses $m_{\\AcB}$ of the two orbiting bodies. To test these predictions, bands are drawn in the $m_{\\ssA} - m_{\\ssB}$ plane, of the form \\be \\xi^{\\rm th}(m_{\\ssA},m_{\\ssB}) = \\xi^{\\rm obs} \\pm \\Delta \\,, \\ee where $\\xi^{\\rm th}$ denotes the theoretically predicted value of a PK parameter, $\\xi^{\\rm obs}$ denotes its observed value and $\\Delta$ denotes the observational error. The validity of GR requires all such bands to have non-empty intersection, providing a substantive test provided at least three PK parameters can be measured. The area of mutual overlap then gives the GR-inferred masses of the stars, to within some tolerance.  Thus far, general relativity has passed these stringent tests, using a number of binary pulsars. For two of these --- PSR B1534+12 \\cite{1534,1534spin} and the double binary, PSR J0737-3039A/B \\cite{wexkramer,dblpsr,dblpsrspin} --- the tests are particularly redundant since it is possible to infer observational values for five independent PK parameters. It is also possible to infer the spin-orbit precession frequency \\cite{1534spin,dblpsrspin}, which can be considered as a sixth PK parameter. Moreover, for the double binary, PSR J0737-3039A/B, both stars are pulsars, and a measurement of the ratio of the semi-major axes yields another constraint on the masses. The masses are constrained quite accurately \\be m_{\\ssA} = 1.3381(7) \\, m_{\\odot} \\quad \\hbox{ and } \\quad m_{\\ssB} = 1.2489(7) \\, m_{\\odot} \\quad \\hbox{(GR)}\\,. \\ee  The success of GR in describing pulsar orbits also constrains alternative theories of gravity; requiring their predictions to agree with GR to within present errors. Prominent among these alternatives are scalar-tensor theories \\cite{sctensrev,defrev}, for which long-range gravitational forces are mediated by both the metric, $g_{\\mu\\nu}$, and a very light scalar, $\\phi$, described by the action \\be S = - \\int \\frac{\\exd^4 x}{c} \\sqrt{-g} \\; \\left[ \\frac{1}{2 \\kappa^2} \\, g^{\\mu\\nu} \\left( R_{\\mu\\nu} + \\partial_\\mu \\phi \\, \\partial_\\nu \\phi \\right) \\right] + S_m \\,. \\ee Here $\\kappa^2 = 8 \\pi G/c^{4} = \\hbar / (M_p^{2} c^{3})$ denotes the gravitational coupling, while $S_m = S_m[g_{\\mu\\nu}, \\phi, \\Psi_i]$ denotes the matter action, and controls how $\\phi$ and $g_{\\mu\\nu}$ couple to the various other `matter' fields, $\\Psi_i$, which we take in what follows to have the form \\be \\label{Smcoup} S_m = S_m[ A^{2}(\\phi) \\, g_{\\mu\\nu}, \\Psi_i] \\,. \\ee This form of coupling has several motivations. First, it is favored by strong observational constraints \\cite{EqPrViolations,TEGP} on violations of the weak equivalence principle, which are evaded by actions of the form of eq.~\\pref{Smcoup}. Second, it is also the kind of theory that actually describes the low-energy limit of certain types of extra-dimensional theories \\cite{XDscalars}.  The predictions of theories of this type have have been compared in detail with binary pulsar data \\cite{sttest, Damour:2007uf}, with the coupling function assumed to have the particular {\\em quasi-Brans Dicke} form, \\footnote{Brans Dicke theory corresponds to the specific choice $b_s = 0$. Our notation follows that of \\cite{HB}, and differs in minor ways from that used elsewhere in the literature \\cite{defrev}, as described in detail in table \\ref{tblnot}.} \\be A(\\phi) = \\exp \\left[ a_{s} \\, \\phi  + \\frac{b_{s}}{2} \\, \\phi^{2} \\right] \\,, \\ee for which the effective coupling of scalars to matter turns out to have the strength \\be \\label{smcoupl} a(\\phi) := \\frac{\\exd \\ln A }{ \\exd \\phi}  = a_s + b_s \\, \\phi \\,. \\ee This form is motivated by the idea that the field $\\phi$ does not vary appreciably in any particular system of interest, and so $a(\\phi)$ is approximately constant.  \\begin{table}[ht] \\begin{center} \\begin{tabular}{c|c|c} {\\bf Our Notation} & {\\bf Notation of \\cite{defrev}} & {\\bf Meaning} \\\\ \\hline\\hline $a_{s}$         & $\\alpha$              & Microscopic scalar-matter coupling constant \\\\ $b_{s}$         & $\\beta$               & Microscopic scalar-matter coupling constant \\\\ $a(\\phi)$       & $\\alpha(\\phi)$        & Microscopic scalar-matter coupling function \\\\ $a_{\\AcB}$      & $\\alpha_{\\AcB}$       & Effective scalar-matter coupling of a star \\\\ $b_{\\AcB}$      & $\\beta_{\\AcB}$        & Effective scalar-matter coupling of a star \\\\ $k_{\\AcB}$      & $\\kappa_{\\AcB}$       & Sensitivity of a star's moment of inertia to scalar field \\\\ $Q_{\\AcB}$      & $\\omega_{\\AcB}$       & Scalar charge of a star \\\\ \\end{tabular} \\end{center} \\caption{\\label{tblnot} Table of Notation.} \\end{table}  Comparison with solar system and pulsar data is found to constrain the microscopic couplings for this theory, $a_s$ and $b_s$, to be consistent with zero, with $a_s$ strongly constrained from solar-system data and $b_s$ restricted by pulsars. The constraints on these microscopic couplings are usually presented as an exclusion plot in the $b_{s}$-$a_{s}$ plane, and in particular it was found that \\cite{sttest, Damour:2007uf}: \\be \\label{bsbound} b_{s} \\gtrsim -4.5 \\,, \\qquad |a_{s} + b_{s} \\phi_{\\infty}^{SS}| < 3.4 \\cdot 10^{-3} \\,, \\ee where $\\phi_{\\infty}^{SS}$ is the value of the scalar field asymptotically far away from the solar system, which is conventionally taken to be zero. \\subsection{A more model-independent approach}  There are two related drawbacks to traditional comparisons between scalar-tensor theory and observations. First, because limits like eq.~\\pref{bsbound} are quoted directly for the microscopic couplings, $a_s$ and $b_s$, a completely new analysis is required for each new assumed functional form for the coupling function $A(\\phi)$. Second, as discussed below (and remarked on by the original authors), these bounds are subject to uncertainties that are hard to quantify, due to limits of our understanding of the nuclear equation of state that applies within the pulsars.  One approach towards robustness that has been taken in the literature is to generalize the form taken for the phenomenological lagrangian describing the two-body interactions of the gravitating objects, restricting attention to theories of gravity whose predictions for the orbital dynamics may be derived from a boost-invariant Lagrangian (at least to the first post-Newtonian (1PN) order). Will and Damour \\& Taylor have shown that such theories may be characterized by a set of body-dependent phenomenological parameters \\cite{TEGP,Damour:1991rd}, of which there are five --- $m_{\\ssA}$, $m_{\\ssB}$, $\\mathcal{G}$, $\\epsilon$, $\\xi$ --- in the most general Lagrangian \\cite{TEGP,Damour:1991rd}: \\be \\mathcal{L} = \\frac{V^{2}}2 + \\frac{\\mathcal{G}M}{R} + \\frac{1}{8c^{2}}(1-3\\nu)V^{4} + \\frac{\\mathcal{G} M }{2R c^{2}} \\left( (\\epsilon + \\nu)V^{2} + \\nu(\\mathbf{N} \\cdot \\mathbf{V})^{2} - \\xi \\frac{\\mathcal{G}M}{R}\\right) \\,, \\ee where $\\mathbf{R}$ and $\\mathbf{V}$ are the relative position and velocity vectors, $\\mathbf{N} = \\mathbf{R}/R$, and $\\mathcal{G}$, $m_{\\ssA}$, $m_{\\ssB}$, $\\epsilon$, and $\\xi$ are phenomenological parameters, and $M=m_{\\ssA} + m_{\\ssB}$, $\\nu = m_{\\ssA}m_{\\ssB}/M^{2}$. Wex and Kramer have used the observed values of the PK parameters for the double pulsar to constrain these WDT (Will-Damour-Taylor) parameters \\cite{wexkramer}.  However, a drawback of this approach is that it does not describe radiation effects, and in particular the PK parameter $\\dot{P}_{b}$ (which describes the shortening of the orbital period due to emission of gravitational radiation). This is because $\\dot{P}_{b}$ cannot be expressed in terms of the WDT parameters. This limitation exists because radiative effects come in at higher order in powers of $V/c$ (at 1.5PN order for dipole emission and 2.5PN order for quadrupole emission), whereas the WDT parameters only describe 1PN effects. Extending the phenomenological parametrization to higher PN orders introduces too many new parameters for the data to usefully constrain, however. Thus, in order to use all PK parameters, including $\\dot{P}_{b}$, to obtain interesting constraints on alternative theories of gravity, it is necessary to further restrict the class of theories that one considers.  In the present paper we take a complementary approach to confronting pulsar observations with scalar-tensor models, based on the observation that the scalar-tensor predictions for the PK parameters depend only on a relatively small number of macroscopic quantities that characterize how the two stars couple to the scalar field. There turn out to be seven of these, (defined in detail in section \\ref{pkpars}): $m_{\\AcB}$, $a_{\\AcB}$, $b_{\\AcB}$ and $k_\\ssA$. These generalize the two masses --- $m_\\ssA$ and $m_\\ssB$ --- that suffice to make predictions within General Relativity, to include four new quantities --- $a_\\AcB$ and $b_\\AcB$ --- that characterize the strength with which the scalar field couples to the pulsar and its orbital partner, plus one variable, $k_\\ssA$, related to the pulsar's moment of inertia.  The key point is that model-dependent complications, like the detailed form of $A(\\phi)$ and knowledge of the nuclear equation of state, enter only into the predictions for the quantities $m_{\\AcB}$, $a_{\\AcB}$, $b_{\\AcB}$ and $k_\\ssA$ as functions of the microscopic couplings in a particular scalar-tensor theory. But these complications do {\\em not} enter at all into the formulae that express how $m_{\\AcB}$ through $k_\\ssA$ determine the observed PK parameters.  This makes it useful to phrase the confrontation between theory and experiment in two steps: first use the observations to constrain the quantities $m_\\AcB$ through $k_\\ssA$ once and for all in a model-independent way; then compute these quantities within specific scalar-tensor models as functions of the underlying model parameters that define the function $A(\\phi)$.  It is the goal of this paper to perform the first --- and model-independent --- one of these steps. At first sight this might seem to be impossible to do, since it appears to involve constraining more quantities than the five observable PK parameters\\footnote{Since the theoretical prediction for the spin-orbit precession frequency has not yet been calculated in scalar-tensor gravity, we do not include it in the present analysis.}. However, it turns out that the present data nonetheless allow useful constraints to be achieved, for two reasons. First, constraints are possible because the dependence of the PK parameters on two of the parameters, $b_{\\AcB}$, is very weak; thus effectively reducing the number of free variables from seven down to five.  Second, one combination of Keplerian parameters --- the ratio of projected semi-major axes of the two orbits, $R \\equiv x_{\\ssB}/x_{\\ssA}$ (see section \\ref{apdx} for details of notation) --- can sometimes also be measured. For instance for the double pulsar observations give $R = 1.0714 (11)$. This measurement is useful because the theoretical prediction for this quantity in all Lorentz-invariant theories of gravity is $R_{\\rm th} = m_{\\ssA}/m_{\\ssB} + \\mathcal{O}(1/c^{4})$ and so is completely independent of the scalar couplings $a_\\AcB$ and $k_\\ssA$ \\cite{Damour:2007uf} (see also section \\ref{apdx}). When this ratio is measurable the number of observational constraints to be satisfied rises from five to six.  In the end we find that interesting model-independent constraints on $m_{\\AcB}$ and $a_{\\AcB}$ are possible. The great virtue of these bounds is that they directly constrain the stellar parameters on which the PK parameters depend, independent of any assumptions about the function $A(\\phi)$ and the nuclear equation of state.  Once the observational constraints on these quantities are known, they can be used to constrain the microscopic parameters for any choice of $A(\\phi)$ or equation of state as a separate step. When we do so for the quadratic model $A(\\phi) = \\exp(a_{s} \\phi + b_{s}\\phi^{2}/2)$, we find agreement with earlier work, which already rules out a large part of the most interesting region (that of spontaneous scalarization \\cite{sttest,Damour:2007uf,spsc}).  It might come as a surprise to the reader that this second step might depend on the details of the neutron equation of state, since it is an important feature of general relativity that predictions for the PK parameters depend only on the two masses, $m_{\\AcB}$, of the stars, and on none of their other properties. This happy property is called the `principle of effacement' of internal structure \\cite{effacement}, and it is this feature that allows a precise prediction of all PK parameters in GR using only the masses, without need for detailed knowledge about the star's structure.  In scalar-tensor gravity (as in most other alternative theories of gravity) the principle of effacement does not hold. The prediction for the PK parameters actually depends in principle on all seven of the quantities which characterize the internal gravitational fields: $m_{\\AcB}$, $a_{\\AcB}$, $b_{\\AcB}$, $k_{\\ssA}$. For now, it is important to note that these are all a-priori independent. Once a particular equation of state is specified, then the equations of stellar structure can be solved, and on a given branch\\footnote{In scalar-tensor theories there may be multiple branches of stellar configurations \\cite{spsc}.} of stellar configurations, the quantities $a$, $b$ and $k$ for each star can be expressed in terms of its mass $m$ and the underlying parameters --- like $a_s$ and $b_s$ --- that define the scalar-tensor model. This is the approach taken in \\cite{sttest, Damour:2007uf}. The resulting bounds are then subject to uncertainties in the equation of state, which are hard to quantify.  In the end, the phenomenological analysis to which we are led in this paper is similar in spirit to that carried out by Wex and Kramer \\cite{wexkramer}. The difference is that we specialize to a more restricted class of theories of gravity, for which the energy loss due to emission of gravitational radiation has been calculated and takes a relatively simple form. Consequently, unlike Wex and Kramer, we are able to make use of all observed PK parameters, including $\\dot{P}_{b}$, when obtaining constraints.  We thus urge observers to express their results in this more model-independent way, which potentially can then be used by theorists to constrain a great variety of specific models. Its independence of the internal structure of the objects involved also shows that some of their physical properties -- like the masses $m_{\\AcB}$ -- can be inferred quite robustly, without making assumptions about which theory of gravity actually applies in Nature. ",
        "conclusions": "In summary, we have shown that the existing data for two pulsars is constraining enough to place model-independent bounds directly on the stellar parameters that control the size of post-Keplerian effects in scalar-tensor models. The virtue of these bounds is that they do not depend on the particular form for the function $A(\\phi)$ that defines which model is of interest, or on the details of the stellar equations of state. In particular we find that existing data impose strong and model-independent constraints on the relative size of the two scalar charges and masses.  These bounds can also be used to constrain particular models by computing in these models the masses and couplings as functions of the microscopic parameters. It is at this point that dependence on things like the stellar equation of state enters.  We applied these methods in particular to the double pulsar J0737-3039A/B, and found that $1.28 \\leq m_{\\ssA}/m_{\\odot} \\leq 1.34$, $1.19 \\leq m_{\\ssB}/m_{\\odot} \\leq 1.25$, $|a_{\\AcB}| < 0.21$, and $|a_{\\ssB}-a_{\\ssA}|< 0.002$, with 68\\% confidence, for all choices of scalar-matter coupling function, and for all nuclear equations of state.  A similar analysis of the pulsar PSR B1534+12 yields $0.97 \\leq m_{\\ssA}/m_{\\odot} \\leq 1.28$, $1.15 \\leq m_{\\ssB}/m_{\\odot} \\leq 1.31$, $|a_{\\AcB}| < 0.44$, and $|a_{\\ssB}-a_{\\ssA}|< 0.004$."
    },
    "1107/1107.3120_arXiv.txt": {
        "abstract": "We present 1.1 mm observations for a sample of 16 powerful radio galaxies at 0.5$<$z$<$5.2 and a radio quiet quasar at z=6.3, obtained using the AzTEC bolometer array mounted on the ASTE or the JCMT.  This paper more than doubles the number of high-z radio galaxies imaged at millimetre/sub-millimetre wavelengths.  We detect probable millimetre-wave counterparts for 11 of the active galaxies.  The 6 active galaxies which do not have a probable millimetre counterpart in our images nevertheless have one or more likely associated millimetric source.  Thus, we conclude that powerful (radio-loud) active galaxies at high-z are beacons for finding luminous millimetre/sub-millimetre galaxies at high-z.  The flux densities of our AzTEC {\\it counterparts} imply star formation rates ranging from $<$200 to $\\sim$1300 $M_{\\odot}$ yr$^{-1}$.  In addition, we find that for the radio galaxoes the 1.1 mm flux density is anticorrelated with the largest angular size of the radio source.  We also present new {\\it Spitzer} imaging observations of several active galaxies in our sample.  Combining these with archival data, we examine the mid-infrared colours of our sample.  We find that radio galaxies for which we have detected a probable 1.1 mm counterpart have mid-infrared colours consistent with dusty starbursts, and are usually bluer than high-z {\\it Spitzer}-selected active galaxies.  In addition, we find arcs of 24 $\\mu$m sources extending across $\\sim$200-500 kpc, apparently associated with three of the radio galaxies. ",
        "introduction": "Powerful, high-z radio galaxies (z$\\ge$0.5: HzRGs) continue to play a key role in cosmological investigations.  Their high luminosities at radio and optical wavelengths make them useful as beacons for finding massive elliptical galaxies and their progenitors out to high redshifts (e.g. Dunlop et al. 1996; R\\\"ottgering et al. 1997; McLure et al. 1999), and provide unique opportunities to study their host galaxy and environment (e.g. Venemans et al. 2007).  Many HzRGs are embedded within haloes of ionized gas which are extended on spatial scales of tens of kpc, and which emit luminous emission lines from various species (e.g. McCarthy, Spinrad \\& van Breugel 1995; Reuland et al. 2003a; Villar-Mart{\\'{\\i}}n et al. 2003).  Spatially resolved kinematic studies of these haloes suggest that the cold gas comprising many of these haloes is in infall towards the centre of the host galaxy (Humphrey et al. 2007), while in others it is being strongly disturbed by the radio jets (e.g. Best, R\\\"ottgering \\& Longair 2000; Nesvadba et al. 2006).  In most cases, ratios between emission lines imply that the radiation field of the AGN is the dominant ionization mechanism for the haloes, although shocks may make a fractional contribution in some cases (e.g. Vernet et al. 2001; Humphrey et al. 2008).  Modelling of the flux ratio Ly$\\alpha$/HeII$\\lambda$1640 measured from one-dimensional spectra has revealed an excess of Ly$\\alpha$ emission above the predictions of AGN photoionization models (Villar-Mart{\\'{\\i}}n et al. 2007), which suggests the presence of star forming regions in the giant gaseous haloes (see also Hatch et al. 2009).  Millimetre emission provides an alternative means by which the star formation history of HzRGs can be probed.  Several tens of HzRGs have now been detected at millimetre wavelengths, with the detection rate rising from $\\sim$15 per cent at z$<$2.5 to $\\ga$75 per cent at z$>$2.5 (Archibald et al. 2001; Reuland et al. 2003).  Millimetric observations are sensitive to cold dust that is re-radiating emission received from young stars: implied rates of star formation range from a few hundred to $\\sim$1500 $M_{\\odot}$ yr$^{-1}$ (Dunlop et al. 1994; Archibald et al. 2001; Reuland et al. 2004).  These millimetre/sub-millimetre detected radio galaxies are not typical of the wider population of millimetre/sub-millimetre selected galaxies (MMGs hereinafter): unlike the vast majority of MMGs, the mm emission from radio galaxies is often extended on spatial scales of $\\ga$100 kpc (e.g. Stevens et al. 2003), can show morphological complexity (Greve et al. 2007), and is not always spatially coincident with the radio-optical nucleus of the AGN and host galaxy (Ivison et al. 2008).  Several HzRGs also appear to have MMGs as companions, or are associated with an overdensity of MMGs (Ivison et al. 2000; Stevens et al. 2003; see also Priddey, Ivison \\& Isaak 2008 and Stevens et al. 2010).  In this paper we present images and measure 1.1 mm flux densities for a sample of 16 powerful radio galaxies at 0.5$<$z$<$5.2 and a radio-quiet quasar at z=6.3, observed with the AzTEC bolometer array instrument (Wilson et al. 2008) mounted on the Atacama Submillimetre Telescope Experiment (ASTE: Ezawa et al. 2004) or the James Clerk Maxwell Telescope (JCMT), as part of the AzTEC/ASTE Cluster Environment Survey (ACES: Zeballos et al., in preparation).  Throughout this paper we adopt a flat Universe with $H_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.27 and $\\Omega_{\\Lambda}$=0.73.  In this cosmology 1\\arcsec corresponds to 5.7-8.6 kpc for 0.5$<$z$<$6.3.   \\begin{table*} \\centering \\caption{Basic properties of the sample of radio galaxies.  Columns are as follow.  (1) Name of the active galaxy.  (2) Active galaxy redshift.  (3) Rest-frame wavelength, in units of $\\mu$m, corresponding to 1.1 mm in the observer frame.  (4) and (5) are the RA and Dec of the active galaxy.  (6) The wavelength regime from which the position of the active galaxy was determined, where R = radio, I = mid-IR (3.6 and 4.5 $\\mu$m), O = optical (R-band) and K = near-IR (K-band).  (7) 1.1 mm flux density of the millimetre counterpart to the active galaxy.  Errors in $S_{1.1}$ are 1$\\sigma$, and upper limits are 3$\\sigma$.  (8) The observed offset of the mm counterpart relative to the position of the active galaxy.  In the absence of a probable counterpart we show in parentheses the offset to the closest mm detection.  (9) P-statistic of the probable mm counterpart.  Again, in the absence of a probable counterpart we show in parentheses the P-statistic of the nearest mm detection.  (10) Log of the far-IR luminosity of mm counterparts.  For non-detections, we give 3$\\sigma$ upper limits.  (11) The star formation rate implied by the far-IR luminosity.  Included at the bottom of the table is the radio-quiet quasar SDSS J1030+0524.} \\begin{tabular}{lllllllllll} \\hline Source name   & z    & $\\lambda_{rest}$ & RA (J2000) & Dec (J2000) & ID & $S_{1.1}$ & r & P & Log $(L_{FIR}/L_{\\odot})$ & SFR \\\\ &      & ($\\mu$m)         &(hh:mm:ss.ss)&(dd:mm:ss.s)&    & (mJy)   &(\\arcsec)& &            & ($M_{\\odot}$ yr$^{-1}$) \\\\ (1)           & (2)  & (3)              & (4)        & (5)         & (6)& (7)       & (8)   &(9)& (10)     & (11)  \\\\ \\hline MRC 2201-555  & 0.51 & 730& 22:05:04.83& -55:17:44.0& I& 6.1$\\pm$0.8 & 2.5    & 8$\\times$10$^{-6}$ & 12.86    & 720    \\\\ MRC 2008-068  & 0.55 & 710& 20:11:14.22& -06:44:03.6& R& 8.6$\\pm$0.9 & 3.5    & 4$\\times$10$^{-6}$ & 13.04    & 1100   \\\\ MRC 2322-052  & 1.19 & 500& 23:25:19.62& -04:57:36.6& I& $<$2.1      & (20.7) & (0.049)            & $<$12.60 & $<$400 \\\\ TXS 2322-040  & 1.51 & 440& 23:25:10.23& -03:44:46.7& R& 2.3$\\pm$0.6 & 2.3    & 0.0006             & 12.68    & 480    \\\\ MRC 2048-272  & 2.06 & 360& 20:51:03.49& -27:03:03.7& I& 2.3$\\pm$0.8 & 11.4   & 0.02               & 12.68    & 480    \\\\ MRC 0355-037  & 2.15 & 350& 03:57:48.06& -03:34:09.5& I& 3.5$\\pm$0.7 & 6.6    & 0.002              & 13.02    & 1040   \\\\ PKS 1138-262  & 2.16 & 350& 11:40:48.35& -26:29:08.6& R& $<$3.6      & (14.3) & (0.007)            & $<$12.85 & $<$710 \\\\ 4C +23.56     & 2.48 & 320& 21:07:14.82& +23:31:45.1& R& $<$1.5      & (21.4) & (0.005)            & $<$12.46 & $<$290 \\\\ MRC 2104-242  & 2.49 & 320& 21:06:58.27& -24:05:09.1& R& 3.7$\\pm$0.9 & 6.9    & 0.002              & 12.87    & 740    \\\\ PKS 0529-549  & 2.58 & 310& 05:30:25.43& -54:54:23.3& I& 6.4$\\pm$0.7 & 3.7    & 2$\\times$10$^{-5}$ & 13.10    & 1270   \\\\ MRC 0316-257  & 3.13 & 270& 03:18:12.14& -25:35:10.2& I& $<$2.1      & (17.3) & (0.009)            & $<$12.55 & $<$350 \\\\ TN J2009-3040 & 3.16 & 260& 20:09:48.08& -30:40:07.4& K& 3.3$\\pm$0.9 & 12.9   & 0.005              & 12.79    & 610    \\\\ 4C +41.17     & 3.79 & 230& 06:50:52.35& +41:30:31.4& O& 3.5$\\pm$1.0 & 6.6    & 0.001              & 12.78    & 610    \\\\ TN J2007-1316 & 3.83 & 230& 20:07:53.22& -13:16:43.4& R& 2.8$\\pm$0.9 & 4.0    & 0.002              & 12.68    & 480    \\\\ TN J1338-1942 & 4.10 & 220& 13:38:26.23& -19:42:33.6& R& $<$5.0      & (15.7) & (0.0008)           & $<$12.90 & $<$790 \\\\ TN J0924-2201 & 5.19 & 180& 09:24:19.91& -22:01:41.5& I& 3.6$\\pm$0.9 & 13.9   & 0.006              & 12.74    & 560    \\\\ \\hline SDSS J1030+0524&6.28 & 150 & 10:30:27.10& +05:24:55.0& I& $<$1.4     & (22.7) & (0.02)             & $<$12.30 & $<$200 \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "In the process of this study, we have identified a statistically significant anti-correlation between the size of the radio source and the brightness of the millimetre continuum emission (or SFR).  This is by no means the first investigation into the possible relationships between the size of a radio source and the other properties of radio galaxies; numerous relationships have been found previously between the observed size of powerful radio sources and various UV, optical and infrared properties.  Best, Longair \\& R\\\"ottgering (1996) found a morphological evolution with radio size in their sample of eight 3CR radio galaxies at 1$<$z$<$1.3.  For galaxies with relatively smaller radio sources, their {\\it HST} images reveal a string of bright knots tightly aligned with the radio emission.  The galaxies with relatively larger radio sources, while still showing the alignment effect (e.g. Chambers, Miley \\& van Breugel 1987; McCarthy et al. 1987), are generally more compact and contain fewer bright optical knots.  Furthermore, galaxies with smaller radio sources, the extended emission line region (EELR) tend to have a relatively smaller observed spatial extent, as measured with Ly$\\alpha$ (van Ojik et al. 1997).  The EELR of these galaxies also show more extreme emission line kinematics, have brighter UV-optical emission lines, and ratios between emission lines that imply they have lower ionization states (Best, R\\\"ottgering \\& Longair 2000; Inskip et al. 2002; Humphrey et al. 2006).  In addition, the polarization of the rest-frame UV continuum emission is correlated with radio source size (Sol\\'orzano-I\\~narrea et al. 2004), suggesting that the ratio between scattered light and starlight is relatively higher in radio galaxies with larger radio sources.  The anticorrelation between 850 $\\mu$m luminosity and UV continuum polarization found by Reuland et al. (2004) appears to confirm that the UV continuum polarization is indeed strongly dependent on the rate of star formation.  Furthermore, HzRGs showing anomalously high values of Ly$\\alpha$/HeII ($\\ga$15), likely due to the presence of a powerful starburst ($\\sim$100 $M_{\\odot}$ yr$^{-1}$) tend to have smaller radio sources than those with 'normal' or 'low' values (Villar-Mart\\'{i}n et al. 2007).  These radio-optical trends are commonly supposed to be the result of (a) interaction between the radio source and the host galaxy, (b) slowing of radio source expansion in radio galaxies that are in relatively denser environments, or (c) age effects of the radio source.  A previously unanswered question was whether the star formation rate is intrinsically higher in HzRGs with smaller radio sources, or whether instead the UV-optically derived SFR merely appears higher due to some difference in dustiness, geometry, etc.  Observations in the millimetre/sub-mm regime give us a new perspective from which we may address this question, since the observations are sensitive to UV-optically obscured (rather than unobscured) star forming regions.  Assuming that the millimetre/sub-mm continuum is dominated by thermal emission from starburst-heated dust, as we believe to be the case, then our anticorrelation between $S_{1.1}$ and radio size leads us to conclude that HzRGs with relatively smaller radio sources do indeed tend to have intrinsically higher SFR.  \\begin{figure} \\special{psfile=radio_z.ps hoffset=-35. voffset=-353. hscale=50. vscale=50.} \\vspace{2.35in} \\caption{Radio size as a function of z.  In the interest of clarity, error bars are not shown.  Note the clear anticorrelation between these two parameters, which complicates interpretation of any observed evolution of properties with z (see text). } \\end{figure}  In light of this conclusion, it seems natural to now question whether the increase in the mm/sub-mm detectability of HzRGs at z$\\ge$2.5 is due to a genuine redshift evolution in the rate or mode of star formation in the host galaxies, as has commonly been supposed (Archibald et al. 2001; Reuland et al. 2004).  Indeed, we propose an alternative hypothesis: that there is a bias towards identifying HzRGs with smaller radio sources at higher z (Fig. 12) because these sources are associated with more luminous narrow emission line regions (e.g. Best, R\\\"ottgering \\& Longair 2000).  And since smaller radio sources typically have more luminous mm/sub-mm emission (Figs. 5 \\& 6), we propose that this translates into a bias in favour of identifying HzRGs with significantly more luminous mm/sub-mm emission at higher z.  Identifying and studying the `missing', line-weak HzRGs will be crucial for testing this hypothesis (see Reuland et al. 2003b; Reuland 2005; Miley \\& De Breuck 2008).  It is not known whether the observed size of the radio source is determined primarily by the density of the interstellar/intergalactic medium through which the source propagates, or whether the size instead determined primarily by the age of the radio source.  If we were to assume the latter, then we would be able to gauge the lifetime of luminous, dusty star formation episodes (S$_{1.1}>$2.5 mJy; SFR$>$500 M$_{\\odot}$ yr$^{-1}$) closely associated with HzRGs.  In Fig. 5, all of the luminous millimetre counterparts are associated with HzRGs that have radio source diameters of d$<$200 kpc: we adopt d$\\sim$200 kpc as the approximate stage in the growth of the radio source at which the luminous millimetre counterparts fade to S$_{1.1}<$2.5 mJy (SFR$<$500 M$_{\\odot}$ yr$^{-1}$).  Assuming a conservative expansion speed of $\\sim$0.01c for the radio source (e.g. Best et al. 1995), this diameter implies an age of $\\sim$30 Myr.  A similar age was estimated by Willott et al. (2002) for a smaller sample of high-z, radio-loud quasars.  However, we must be cautious about generalising this to obscured star formation in non-active galaxies (i.e., to sub-mm galaxies in general): HzRGs contain two powerful sources of feedback, the radio source and the accretion disc, both of which are able to quench star formation and, perhaps, to ignite it.  Concievably, feedback of this kind could either lengthen or shorten the lifetimes of star formation activity in HzRGs, and could mean that lifetimes of obscured star formation activity determined for HzRGs are not necessarily representative of those for millimetre/sub-millimetre galaxies in general.  We have also investigated whether there might be a correlation between the millimetre flux density and the Ly$\\alpha$/HeII ratio, which is expected to be sensitive to SFR also.  While Ly$\\alpha$ can originate from stellar photoionized HII regions and from AGN-ionized nebulae, the HeII line is expected to come predominantly from nebulae that are ionized by the AGN.  Though both Ly$\\alpha$/HeII and the millimetre flux density are expected to be correlated with SFR, and therefore with each other, no significant trend was found between them.  Furthermore, the luminosity of the Ly$\\alpha$ emission is often more than an order of magnitude lower than one would have expected given the SFR derived from the 1.1 millimetre flux density (we estimate $L_{Ly\\alpha}$/$L_{FIR}\\sim$0.0007-0.054, while the expected value is $\\sim$0.03).  This may mean that one (or both) of the millimetre and Ly$\\alpha$ emission is an inaccurate tracer of the SFR in HzRGs.  This could be due to destruction of Ly$\\alpha$ photons by the dust associated with this kind of star formation activity.  Thus far, we have tacitly assumed that the sources we identify as 1.1 mm counterparts are the radio galaxy hosts themselves.  However, it is important to be aware that the spatial resolution of our observations, together with the signal to noise ratios of the detections, has resulted in 1$\\sigma$ positional uncertainties of $\\sim$2\\arcsec-16\\arcsec for our 1.1 mm counterparts.  This means that the sources we identify as 1.1 mm counterparts could in actual fact be offset by tens or even hundreds of kpc from the galactic nucleus, which would place them in the outskirts or outside of the host galaxy.  This may be the case for TN J0924-2201, which is undetected in SCUBA photometry (Reuland et al. 2004), but which according to our criteria has a 1.1 mm counterpart in our AzTEC map ($\\S$3.2) -- this suggests that the AzTEC source is not the actual counterpart, but is instead a MMG several tens of kpc from the radio galaxy.  Also of relevance in this context are results from a FWHM=2\\arcsec spatial resolution SMA study of the z$=$3.79 radio galaxy 4C +60.07 (Ivison et al. 2008): the luminous sub-mm emission, previously assumed to be centred on the active galaxy's host, was resolved into two discrete sub-mm sources with offsets of 1\\arcsec ($\\sim$10 kpc) and 4\\arcsec ($\\sim$30 kpc) from the position of the radio core.  These two results provide evidence, although anecdotal in nature, that the millimetre/sub-mm emission associated with high-z radio galaxies may not always be spatially coincident with the host galaxy.  Unfortunately, with currently available facilities high resolution observations of the kind presented by Ivison et al. (2008) are costly and are sensitive only the very brightest of MMGs.  The next generation of millimetre/sub-mm facilities, such as the 50 m Large Millimetre Telescope or the Atacama Large Millimetre Array, will be needed to resolve this issue.  Nevertheless, we can state that all 17 of the active galaxies in our sample are associated with one or more MMG, and thus we conclude that high-z (radio-loud) active galaxies are beacons for finding MMGs."
    },
    "1107/1107.4911_arXiv.txt": {
        "abstract": "We study the K-Inflation models where the inflaton field has non-canonical kinetic term. In particular, we consider the Dirac-Born-Infeld (DBI) form for the kinetic energy of the inflaton field. We consider quadratic and quartic potentials as well as the potential for the natural inflation. We use a modified version of the MODECODE \\cite{Mortonson:2010er} to calculate the power spectrum of the primordial perturbations generated by the inflaton field and subsequently use the WMAP7 results  to constrain the models. Interestingly with DBI type kinetic term, lesser gravity waves are produced as one approaches more towards scale invariance. This is true for all the potentials considered. Unlike the canonical case, this feature, in particular, helps the quartic  ($\\lambda\\phi^4$) potential  with DBI type kinetic term to be consistent with WMAP data. ",
        "introduction": "In the last two decades, there have been number of breakthroughs in our understanding of the Universe due to some remarkable observational results.  COBE first detected the anisotropy in the cosmic microwave background radiation (CMBR) in 1992 \\cite{Smoot:1992td}; subsequently WMAP \\cite{Spergel:2003cb,Peiris:2003ff,Spergel:2006hy,Komatsu:2008hk,Komatsu:2010fb} measured these anistropies with extraordinary  accuracy  that revolutionized our understanding of the physical Universe (also see the recent results by Atacama Cosmology Telescope (ACT)\\cite{Sudeep2011a,Sudeep2011b}).  This has been equally supplemented by the discovery of dark energy by different groups while studying the light curves of different Type-Ia supernovae at different redshifts \\cite{sn} as well as the measurement of large scale structure by various galaxy-redshift surveys \\cite{lss}. All these observational results strongly support the idea that the early universe had a short period  of accelerated expansion. This epoch of accelerated expansion in the early universe, which is also termed as \"inflation\", \\cite{Guth:1980zm, Linde:1981mu,Starobinsky:1980te} not only solves the horizon, flatness and monopole problems in the standard cosmology, it is also the simplest and most favored paradigm for the generation of the nearly scalar invariant super-horizon fluctations in the early universe  \\cite{Mukhanov:1981xt,Hawking:1982cz,Starobinsky:1982ee,Guth:1985ya}, which is believed to be the origin of CMBR anisotropy and the large scale structures in our present day Universe.  Simplest model of inflation is usually described by a scalar field (termed as inflaton) rolling over its potential ( also see \\cite{Debajyoti:2011} for a different approach for obtaining inflation without scalar fields). The field is minimally coupled to the gravity sector. For a sufficiently flat potential, the potential energy of the field dominates over its kinetic energy which essentially gives rise to negative pressure. This drives the accelerated expansion. This scenario is usually described as ``Slow-Roll Approximation\" for inflation \\cite{slowroll}.  The primordial perturbations which are responsible for the structures in our universe can be traced back to the quantum fluctuations of this inflaton field in the very early universe. Any given scalar field inflationary scenario can give rise to both the scalar and tensor fluctuations which can be tested against observational data from the CMBR anisotropy measured by WMAP or other experiments.  With the quality of present data, the primordial fluctation of the inflation field is essentially described by a set of two parameters: the amplitude $A_{s}$ and the spectral index $n_{s}$ of the power specturm of the scalar mode of the density perturbation. As the data from future experiments, such as Planck \\cite{planck}, improve, one can add parameters related with primordial gravitational waves or nongaussianity to this set.  Recently Mortonson et al. \\cite{Mortonson:2010er}  have solved the inflationary mode equations numerically without the usual slow-roll approximation. For this, they proposed a new code called \"MODECODE\" that has been interfaced with the CAMB \\cite{Lewis:1999bs}  and COSMOMC \\cite{Lewis:2002ah} to compute the CMBR anisotropy power spectra and to perform the MCMC analysis for the parameter estimation.  They subsequently used this code to constrain the single scalar field models having canonical kinetic energy term.  They have used the power  law potentials ($V(\\phi) = \\phi^n$, with $n = 2.3, 1, 2$ and $4$) as well as the natural \\cite{Freese:1990rb, Adams:1992bn} and hilltop potentials \\cite{hilltop} for their study. The study shows that the current CMB data put stringent constraints on the parameters of a number of well motivated and interesting single field inflationary models. As an example, they showed that the WMAP7 data excludes the quartic potential ($V(\\phi) = \\lambda \\phi^4$) model.  There are large number of attempts to embade inflation in some high energy physics model. One of them is k-inflation \\cite{kinf}, where the inflation is achieved by the non-standard kinetic terms of the inflaton. One of the examples of such a non-canonical kinetic term is the Dirac-Born-Infeld (DBI) form \\cite{dbi}. Recently  scalar field with DBI type kinetic term has been widely studied in cosmology for both early and late time accelerated expansion of the Universe \\cite{dbiref}. One popular equation of state that arises due to such a kinetic term is so called Chaplygin gas equation of state \\cite{gcg} which has many interesting cosmological signatures.  In the present investigation, we extend the work by Mortonson et al. for scalar fields with DBI type of kinetic term. For this we modify the MODECODE to include the DBI type kinetic term. We study the power law potentials, ($V(\\phi) = \\phi^n$  with $n =2$ and $4$)  as well as the potential for natural inflation. ",
        "conclusions": "In brief, we study the inflation driven by the scalar field models with non-canonical kinetic term, more specifically kinetic term of DBI form. To calculate the scalar and tensor power spectrum generated by the quantum fluctuations of these fields, we modify the publicly available MODECODE to incorporate the non-canonical kinetic term. We consider the power law as well as the natural inflation potentials for our study.  There are number of interesting results that emerge from our study. Unlike the canonical case,  parameter $N_{pivot}$, is almost uncorrelated with other model parameters. Moreover, the shape of the contours in the $r-n_{s}$ plane for all the potentials show that as one approaches towards scale invariance, lesser gravity waves are produced. This result is completely opposite to that for the canonical case. This helps the $\\lambda\\phi^4$ potential with a non canonical kinetic term to be allowed by the observational data. There is also a substantial improvement in the Maximum Likelihood value for the $\\lambda\\phi^4$ potential whereas for other potentials, we get similar number as obtained by Mortonson et al. \\cite{Mortonson:2010er} for the canonical case.  In our study, we consider the action where the gravity sector is minimally coupled with the scalar field, having non-standard kinectic term.  This can be easily extended to models where the scalar field is nonminimally coupled with the gravity sector. This will be our aim in future."
    },
    "1107/1107.5310_arXiv.txt": {
        "abstract": "We present the second of two papers concerning the origin and evolution of local early-type galaxies exhibiting dust features. We use optical and radio data to examine the nature of active galactic nucleus (AGN) activity in these objects, and compare these with a carefully constructed control sample.  We find that dust lane early-type galaxies are much more likely to host emission-line AGN than the control sample galaxies. Moreover, there is a strong correlation between radio and emission-line AGN activity in dust lane early types, but not the control sample. Dust lane early-type galaxies show the same distribution of AGN properties in rich and poor environments, suggesting a similar triggering mechanism. By contrast, this is not the case for early types with no dust features. These findings strongly suggest that dust lane early-type galaxies are starburst systems formed in gas-rich mergers. Further evidence in support of this scenario is provided by enhanced star formation and black hole accretion rates in these objects. Dust lane early types therefore represent an evolutionary stage between starbursting and quiescent galaxies. In these objects, the AGN has already been triggered but has not as yet completely destroyed the gas reservoir required for star formation. ",
        "introduction": "\\label{sec:introduction}  Dust lane early-type galaxies are those galaxies classified in Galaxy Zoo~2\\footnote{http://zoo2.galaxyzoo.org} as early-type galaxies with dust lane features. The classification procedure and construction of the dust lane sample are described in detail in a companion paper, \\citet[hereafter Paper~I]{kav12}. In that paper, we showed that dust lane early-type galaxies often exhibit morphological disturbances, show signs of recent star formation and have stellar ages that are older than starburst early types, but younger than the overall early-type population. We therefore argued that these objects are starburst systems, and suggested that they might arise from gas-rich minor mergers. While intergalactic medium accretion is also a plausible mechanism, we ruled it out in this case due to the presence of large amounts of dust, and the disturbed nature of most galaxies.  In the present paper we test this paradigm by examining the active galactic nucleus (AGN) properties of dust lane early-type galaxies. Numerous theoretical and observational arguments connect star formation and AGN triggering in galaxies. The basic point is that similar conditions (namely the presence of cold gas) are required for both a nuclear starburst and AGN fuelling. Once the supermassive black holes that reside at galaxy centres become active, radiative and kinetic feedback from these objects can profoundly affect their environments through heating and expulsion of gas. AGN feedback has been invoked to explain the cooling flow problem \\citep{bin95}, and accounts for both the shape of bright end of the optical luminosity function \\citep{sha09} and the strong correlation between black hole and bulge mass \\citep{mag98,hae04} in the local Universe.  A number of diagnostics are used to identify AGN. Perhaps the best known is the so-called BPT \\citep*{bal81} analysis, which relies on emission-line measurements to separate star-forming galaxies from AGN. Synchrotron emission from AGN jets and jet-inflated cocoons manifests itself at radio frequencies. X-ray and infrared diagnostics are also used.  Previous studies have shown that Seyfert and low-ionization nuclear emission-line region (LINER) elliptical hosts are more likely to contain dust than inactive ellipticals \\citep{mar03,sim07}. However, it is unclear whether this finding can be extended to other AGN classifications. Recently, \\citet{bes05b} showed that radio and emission-line AGN activities are largely independent at low radio luminosities, in contrast to the well-known correlation at high luminosities \\citep{raw91}. This led \\citet*{har07} to suggest an environmental origin for this dichotomy. Bright radio sources preferentially live in the field; these are also the objects that often exhibit emission-line AGN properties. By contrast, low-luminosity radio AGN tend to live in denser environments, and show little emission-line activity. \\citeauthor{har07} argued that fundamentally different AGN fuelling mechanisms are likely to be at work in these two scenarios, with radio-loud objects in dense environments mainly fuelled by the relatively slow bremsstrahlung cooling out of the hot X-ray gas halo, while in poor environments mergers are primarily responsible for the sudden increase in cold gas content. The properties of the black hole accretion disc are different in the two cases. At high accretion rates associated with mergers, the disc assumes the standard geometrically thin, optically thick shape \\citep[e.g.][]{sha73}; this disc is radiatively efficient and can also produce a radio jet. Thus, both emission-line and radio AGN are detected. At low accretion rates, however, the accretion disc puffs up and becomes radiatively inefficient, yielding a powerful radio jet but no corresponding emission lines \\citep{mei01}. Therefore, a detailed analysis of both optical and radio AGN properties of our sample can help reconstruct the origins of the gas used in both star formation and AGN fuelling, and indicate whether dust lane early-type galaxies are indeed produced in galaxy interactions. Moreover, sizes and luminosities of radio AGN can be used to constrain their ages \\citep{sha08} and therefore place these objects on an evolutionary sequence.  The homogeneity and size of our sample of dust lane early-type galaxies allow us to study the AGN-triggering mechanisms via a combination of photometry, spectroscopy and radio data. We create well-defined control samples, facilitating the construction of an evolutionary sequence and uncovering the significance of the observed dust deatures.  The paper is structured as follows. In Section~\\ref{sec:SampleConstruction} we present our sample of dust lane early-type galaxies as well as a relevant control sample. In particular, we mimic the salient properties of dust lane galaxies in constructing the control sample. Section~\\ref{sec:AGNselection} outlines the AGN selection procedure using two alternative definitions: radio luminosity and emission lines. We present our results in Section~\\ref{sec:results}. Environmental dependence of dust lane AGN properties indicates that the origin of these objects is merger driven. We show that these mergers are necessarily recent, and gas-rich. Finally, we conclude in Section~\\ref{sec:conclusions}.  Throughout the paper, we adopt the Hubble constant of $H_0=72$\\,km\\,s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusions}  We consider the origin and evolution of dust lane early-type galaxies. In a companion paper, we showed that these objects are preferentially located in poor environments, are morphologically disturbed, and have higher star formation rates than the mean for early-type galaxies.  In this work we addressed the issue of the origin of these objects. By employing optical (SDSS) and radio (FIRST and NVSS) surveys we examined AGN properties of dust lane early-type galaxies and a relevant control sample of early-type. Our findings are as follows.  \\begin{enumerate} \\item[(1)] Dust lane early types are much more likely to host emission-line AGN than the control sample of early-type galaxies. \\item[(2)] Radio-loud AGN in dust lane early types do not show an environmental dependence, with similar fractions detected in the field and groups/clusters. By contrast, control sample early-type galaxies are much more likely to be radio loud if located in a cluster. \\item[(3)] Radio-loud AGN in dust lane early types are much more likely to also be detected in emission lines than radio-loud AGN in the control sample. \\end{enumerate}  These results all strongly suggest that minor mergers act as a trigger for AGN activity in dust lane early-type galaxies.  \\begin{enumerate} \\item[(4)] Dust lane early types have starburst ages that are younger than the mean for early-type galaxies. \\item[(5)] Dust lane early-type galaxies exhibit higher star formation and black hole accretion rates than the control sample of early types with the same starburst ages. \\end{enumerate}  These results indicate that the minor mergers are both gas rich and recent. The mergers act as a trigger for both the starburst and AGN activity. The dust lane phase is associated with starburst galaxies in which the AGN has already been triggered, but has not yet shut off star formation completely."
    },
    "1107/1107.0703_arXiv.txt": {
        "abstract": "{Ground-based imaging and imaging spectropolarimetric data are often subjected to post-facto reconstruction techniques to improve the spatial resolution.} {We test the effects of reconstruction techniques on two-dimensional data to determine the best approach to improve our data.}{We obtained an 1-hour time-series of spectropolarimetric data in the \\ion{Fe}{i} line at 630.25\\,nm with the G{\\\"o}ttingen Fabry-P{\\'e}rot Interferometer (FPI) that are accompanied by imaging data in the blue continuum at 431.3\\,nm and \\ion{Ca}{ii} H at 396.85\\,nm. We apply both speckle and (multi-object) multi-frame blind deconvolution ((MO)MFBD) techniques. We use the ``G\\\"ottingen'' speckle and speckle deconvolution codes and the MOMFBD code in the implementation of Van Noort et al.~(2005). We compare the resulting spatial resolution and investigate the impact of the image reconstruction on spectral characteristics of the G\\\"ottingen FPI data.} {The speckle reconstruction and MFBD perform similar for our imaging data with nearly identical intensity contrasts. MFBD provides a better and more homogeneous spatial resolution at the shortest wavelength when applied to a series of image bursts. The MOMFBD and speckle deconvolution of the intensity spectra lead to similar results, but our choice of settings for the MOMFBD yields an intensity contrast smaller by about 2\\,\\% at a comparable spatial resolution. None of the reconstruction techniques introduces significant artifacts in the intensity spectra. The speckle deconvolution (MOMFBD) has a rms noise in Stokes $V/I$ of 0.32\\,\\% (0.20\\,\\%). The deconvolved spectra thus require a high significance threshold of about 1.0\\,\\%  to separate noise peaks from true signal. A comparison to spectra with a significantly higher signal-to-noise (S/N) ratio and to spectra from a magneto-hydrodynamical simulation reveals that the G{\\\"o}ttingen FPI can only detect about 30\\,\\% of the polarization signal present in quiet Sun areas. The distribution of NCP values for the speckle-deconvolved data matches that of observations with higher S/N better than MOMFBD, but shows seemingly artificially sharp boundaries and unexpected changes of the sign.}{For our imaging data, both MFBD and speckle reconstruction are equivalent, with a slightly better and more stable performance of MFBD. For the spectropolarimetric data, the higher intensity contrast of the speckle deconvolution is balanced by the smaller amplification of the noise level in the MOMFBD at a comparable spatial resolution. The noise level prevents the detection of weak and diffuse magnetic fields. Future efforts should be directed to improve the S/N of the G{\\\"o}ttingen  FPI spectra for spectropolarimetric observations to lower the final significance thresholds.} ",
        "introduction": "One of the primary objectives of modern solar physics is the determination of the thermodynamic and magnetic structure of the solar atmosphere at the smallest spatial scales, including isolated magnetic elements in the quiet Sun and the internal fine-structure of larger features such as sunspots. An accurate derivation of the atmospheric properties requires both high spatial and spectral resolution. Telescopes with apertures of 1\\,--\\,1.5\\,m, such as the Swedish Solar Telescope \\citep[SST,][]{scharmer+etal2003}, GREGOR \\citep{volkmeretal07,balthasar+etal2007}, or the New Solar Telescope \\citep[][]{denkeretal06}, allow one to study the dynamics of the solar fine-structure at spatial scales down to 50\\,--\\,100 km on the surface of the Sun. The next-generation solar telescopes of the 4-m class, i.e., the Advanced Technology Solar Telescope \\citep [][]{wagneretal08} or the European Solar Telescope \\citep [][]{collados08}, then will resolve the fundamental scale of the solar photosphere, i.e., the photon free path length.  All observations from the ground are, however, degraded by the wavefront distortions caused by fluctuations of the refractive index in the Earth's atmosphere. To overcome this limitation, adaptive optics (AO) systems are used to correct the wavefront errors, but their compensation is only partial and the diffraction limit of the respective telescope is hard to reach. AO systems have been installed at the Dunn Solar Telescope \\citep[DST,][]{rimmele+radick98, rimmele2004a}, the German Vacuum Tower Telescope \\citep[VTT,][]{vdluehe+etal2003}, and the SST \\citep{scharmer+etal2003a}. AO and multi-conjugated AO systems \\citep{beckers1988,berkefeld+etal2001,moretto+etal2004,ellerbroek+vogel2009} with a larger corrected field of view (FOV) are also foreseen for all next-generation telescopes.  For spectroscopic or spectropolarimetric observations taken with slit-spectrograph systems, it is difficult to improve the spatial resolution beyond the limit provided by the real-time correction of an AO system because of the essentially one-dimensional information at one position of the slit \\citep[see, e.g.,][submitted to A\\&A]{keller+johannesson1995,suetterlin+wiehr2000,aschwanden2010,beck+rezaei2011}. Contrary to that, instruments providing instantly a two-dimensional FOV offer the possibility for a post-facto reconstruction of the observations. Two-dimensional (2D) spectrometers are commonly realized in the form of Fabry-P{\\'e}rot Interferometers (FPIs) that have a narrow tunable spectral band-pass and a high transmission.  A first example of a 2D FPI spectrometer with three etalons was designed for the Australian solar facility in Narrabri \\citep{ramsayetal70, loughheadetal78, bray88}. Nowadays, almost all ground-based solar telescopes are equipped with FPI instruments: the Italian Panoramic Monochromator \\citep{bonaccinietal89,cavallini98} at the THEMIS telescope, the Interferometric BIdimensional Spectrometer \\citep[IBIS,][]{cavallinietal00, cavallini2006, reardon+cavallini2008} at the DST, the visible-light and near-infrared imaging magnetographs \\citep{denkeretal03} at the Big Bear Solar Observatory, the CRisp Imaging Spectro Polarimeter \\citep[CRISP, e.g.,][]{vannoort+vandervoort2008,scharmeretal08} at the SST, and the Triple Etalon SOlar Spectrometer \\citep[TESOS,][]{kentischeretal98, vdluehe+kentischer00, tritschleretal02} at the VTT. For the VTT, another 2D instrument was developed prior to TESOS by the Institute for Astrophysics G\\\"ottingen, the so-called G\\\"ottingen FPI \\citep[][]{bendlinetal92, bendlin93, bendlin+volkmer93, bendlin+volkmer95}. This instrument will in the future be one of the post-focus instruments at the GREGOR telescope. Therefore, the spectrometer was largely renewed during the first half of 2005 \\citep{puschmann+etal2006}. Its design as the new ``Gregor FPI'' is described in \\citet{puschmannetal07} and \\citet{denker+etal2010}.  Even if most of the 2D instruments were initially designed as spectrometers only, they now have all been upgraded for (vector) spectropolarimetry (see e.g., \\citet{nazi+kneer2008a} and \\citet{balthasaretal09} for the G{\\\"o}ttingen FPI; \\citet{viticchi+etal2009} and \\citet{judge+etal2010} for IBIS, and \\citet{beck+etal2010} for TESOS). All 2D instruments use at least two channels, where one beam passes through the tunable FPIs with their high spectral resolution (narrow-band (NB) channel). The second channel uses only a broad-band (BB) interference filter to provide strictly simultaneous images with a high signal-to-noise (S/N) ratio for an accurate alignment of the individual wavelength steps and the application of post-facto reconstruction techniques.  Two main reconstruction techniques have been developed and successfully applied to observation data up to now: the speckle reconstruction technique with the possibility of a subsequent NB deconvolution of spectroscopic or spectropolarimetric data, and the blind deconvolution (BD) technique that can be applied to both imaging and 2D spectropolarimetric data.  For a speckle reconstruction, a fast series of about 50\\,--\\,100 short-exposed images $I_{\\rm obs}({\\bf x},t_i)$ is taken at times $t_i$. The observed images represent the image of the real object $I_{\\rm real}({\\bf x})$ modulated by the instantaneous optical transfer function (OTF) of the Earth's atmosphere and the telescope. In Fourier space, the corresponding equation reads \\begin{equation} \\tilde I_{\\rm obs}({\\bf s}, t_i) = \\tilde I_{\\rm real}({\\bf s}) \\cdot {\\rm OTF} ({\\bf s}, t_i) \\;, \\label{eq_otf} \\end{equation} where $^\\sim$ denotes the Fourier transform, and ${\\bf x}$ and ${\\bf s}$ are the spatial coordinates and their corresponding Fourier equivalent, respectively.  To recover the information of the real image, the Fourier amplitude and phases of $\\tilde I_{\\rm real}$  are derived separately. The Fourier amplitudes are usually determined by the method of \\citet{labeyrie1970} and the spectral ratio method \\citep[][]{vdluehe1984}. Two main approaches exist for the estimate of the Fourier phases, i.e., the Knox-Thompson and speckle masking algorithm \\citep{knox+thompson1974,weigelt1977}. The large number of input images is necessary to obtain good atmospheric statistics. For the application to solar observations with low-contrast images, several specific speckle reconstruction methods were developed \\citep[see, e.g.,][]{vdluehe1993,deboer93,mikurda+vdl2006}. Examples of recent speckle codes are the Kiepenheuer-Institute Speckle Interferometry Package \\citep[KISIP,][]{woeger+etal2008,woeger+vdluehe2008} and the G\\\"ottingen speckle reconstruction code \\citep{deboer93}. The latter was improved by \\citet{puschmann+sailer2006} to account for the field-dependent real-time correction by AO (see also \\citet{denker+etal2007}, or \\citet{woeger2010} for the KISIP code) and successfully applied in, e.g., \\citet{puschmannwiehr06}, \\citet{blancoteal07}, or \\citet{sanchezetal08}. For 2D spectro(polari)meters, the BB images can be directly speckle reconstructed because a sufficient number of images is available. The simultaneously recorded NB images -- always only a few images per wavelength position -- can then be deconvolved using the method of \\citet{keller+vdluehe1992} estimating the instantaneous OTF from the observed and reconstructed BB images \\citep[see also][]{kriegetal99,janssen03,bellogonzalezetal05}.  In the case of the BD technique, one estimates both the real image {\\em and} the OTF at the same time without using an OTF model based on the seeing statistics. This leads to a problem similar to the disentanglement of magnetic field strength and area filling factor for spectral lines in the weak-field limit: an infinite number of possible combinations can yield nearly the same result. The ambiguity is lessened by the use of multiple image frames, if the real image is assumed to be always the same during the observed fast image sequence. Compared with the speckle reconstruction, a significantly smaller number of images, about 5\\,--\\,20, is sufficient \\citep[see, e.g.,][]{loefdahl+etal2007}. Such a multi-frame blind deconvolution \\citep[MFBD,][]{schulz93,vankampen+paxman98,loefdahl2002} can be applied to series of images of a single real object. The approach was extended by \\citet{loefdahl2002} and \\citet{vannortetal05} to cover also the case of multiple objects, e.g., simultaneous images of the same object in slightly different wavelengths or different polarization states. This allows the reconstruction of the simultaneously recorded BB and spectro(polari)metric NB images of, for instance, FPI instruments. The corresponding multi-object multi-frame blind deconvolution \\citep[MOMFBD,][]{vannortetal05} code is freely available\\footnote{http://www.momfbd.org/}. A drawback of the (MO)MFBD is the necessary computing effort that usually requires the use of a dedicated computer cluster or similar installations, whereas speckle codes can also be run within limitations on a single desktop PC or nearly in real-time on a computer cluster \\citep{denker+etal2001}.  Even if the application of reconstruction methods to imaging and 2D spectroscopic or spectropolarimetric data is a routine task these days, there are only relatively few investigations on the influence of the reconstruction on the data properties, e.g., \\citet{mikurda+etal2006} for spectroscopic observations or \\citet{vannoort+vandervoort2008} for MOMFBD-reconstructed data. The reconstruction introduces more severe changes of the data than, e.g., a simple flat fielding, and therefore also may introduce systematic artifacts.  FPI instruments sequentially scan the wavelength points; the {\\em independent} observation and reconstruction of individual wavelength steps thus may compromise the spectral quality. This is even worse for spectropolarimetric observations, where additionally the information on the polarization signal is encoded in the intensity of various modulation states. An independent scaling of the intensity in the individual images -- varying across the FOV as well -- therefore may easily introduce spurious polarization signals because they are finally derived from intensity differences.  In this contribution, we investigate the performance of speckle and BD reconstruction techniques. Because our observations provided us with both imaging and 2D spectropolarimetric data, we extended the study to both types of data. In Sect.~\\ref{sect_obs}, we describe the observational setup and the data acquisition. Section \\ref{sect_datared} explains the data reduction and reconstruction methods for the individual data sets. The results presented in Sect.~\\ref{sect_results} are summarized and discussed in Sect.~\\ref{sect_summ}. Section \\ref{concl} provides the conclusions. ",
        "conclusions": "} To determine the ``best'' approach for the deconvolution of polarimetric spectra, one has to choose a compromise between spatial resolution, intensity contrast, and the noise level in the polarization signal. The destretched spectra provide a reference for the lowest achievable noise level that in our case is roughly preserved by the MOMFBD of the spectra with equal weight for BB and NB channel. For the future usage of the G{\\\"o}ttingen FPI for imaging spectropolarimetry it, however, seems to be necessary to improve the S/N ratio, to be able to lower the final significance thresholds. Especially for weak spectral lines with a small line depth, and hence, small polarization amplitudes, or strong chromospheric lines with a low line-core intensity, the most common polarization amplitudes are far below the 1\\,\\% level and thus cannot be detected at a low S/N."
    },
    "1107/1107.0968_arXiv.txt": {
        "abstract": "Magnetic fields in the intra-cluster medium of galaxy clusters have been studied in the past years through different methods. So far, our understanding of the origin of these magnetic fields, as well as their role in the process of structure formation and their interplay with the other constituents of the intra-cluster medium is still limited. In the next years the up-coming generation of radio telescopes is going to provide new data that have the potential of setting constraints on the properties of magnetic fields in galaxy clusters. \\\\ Here we present zoomed-in simulations for a set of massive galaxy clusters ($M_{v} \\geq 10^{15} h^{-1 }M_{\\odot}$). This is an ideal sample to study the evolution of magnetic field during the process of structure formation in detail. Turbulent motions of the gas within the ICM will manifest themselves in a macroscopic magnetic resistivity $\\eta_m$, which has to be taken explicitly into account, especially at scales below the resolution limit.  We have adapted the MHD {\\tt GADGET} code by Dolag \\& Stasyszyn (2009) to include the treatment of the magnetic resistivity and for the first time we have included non-ideal MHD equations to better follow the evolution of the magnetic field within galaxy clusters. We investigate which value of the magnetic resistivity $\\eta_m$ is required to match the magnetic field profile derived from radio observations. We find that a value of $\\eta_m \\sim 6 \\times 10^{27}$ cm$^2$s$^{-1}$ is necessary to recover the shape of the magnetic field profile inferred from radio observations of the Coma cluster. This value agrees well with the expected level of turbulent motions within the ICM at our resolution limit.  The magnetic field profiles of the simulated clusters can be fitted by a $\\beta-$model like profile (Cavaliere \\& Fusco-Femiano 1976), with small dispersion of the parameters.  We find also that that the temperature, density and entropy profiles of the clusters depend on the magnetic resistivity constant, having flatter profiles in the inner regions when the magnetic resistivity increases. ",
        "introduction": "\\label{sec:intro} Magnetic fields are an important ingredient to understand the physical processes taking places in the intra-cluster medium (ICM) of galaxy clusters. Their presence is demonstrated by radio observations, which, since the last 30 years, have revealed diffuse and faint radio sources filling the central Mpc$^3$ of some galaxy clusters \\citep[radio halos, see \\eg][]{2009A&A...507.1257G,2008A&A...484..327V}. These sources arise because of the interaction of highly relativistic electrons with the ICM magnetic fields.  About 30 radio halos are known so far, and all of them are found in clusters with clear signatures of on-going or recent merger activity \\citep[\\eg][]{2001ApJ...553L..15B,2001A&A...369..441G, 2010ApJ...721L..82C}. The origin of the relativistic particles still needs to be understood, although several models have been proposed. Shocks and turbulence associated with merger events are expected to inject a considerable amount of energy in the ICM, that could compress and amplify the magnetic field and (re-)accelerate relativistic electrons, giving thus rise to the observed radio emission \\citep[see][ for recent reviews of the subject]{2008SSRv..134...93F,2008SSRv..134..311D}. Understanding the magnetic field amplification and evolution during the process of structure formation is then mandatory for modeling the acceleration, transport and interactions of non-thermal energetic particles and thus to understand the observed emission. In addition, an accurate modeling of the magnetic field properties is necessary to understand both the heat transport and the dissipative processes in the ICM.\\\\ The properties of magnetic fields in the ICM have been investigated in the past through cosmological simulations, performed with different numerical codes \\citep{1999A&A...348..351D,2002A&A...387..383D,2005JCAP...01..009D, 2008A&A...482L..13D, 2009MNRAS.398.1678D, 2010ApJS..186..308C} and also through Faraday Rotation measures analysis \\citep[\\eg][]{2004A&A...424..429M, 2006A&A...460..425G, 2005A&A...434...67V, 2008MNRAS.391..521L, 2010A&A...513A..30B}. The comparison with observed data is necessary to constrain the main magnetic field properties, and it is starting now to be feasible thanks to the progress that has been done in the recent years. One key aspect is that, so far, large scale radio emission is mainly detected in very massive clusters. Such massive systems are not easily studied by numerical simulations, since the size of the density fluctuations responsible for the formation of massive halos is large, \\iee  $\\sim$ 20 $h^{-1}$Mpc, and the value of the cosmological parameter $\\sigma_{8}$ in the standard $\\Lambda CDM$ model requires that statistically a total volume of $\\sim 200 h^{-1} Mpc^3$ needs to be sampled by simulations in order to produce at least one cluster as massive as $\\sim 10^{15} h^{-1}$\\msun.  An important step for studying non-thermal phenomena is to perform simulations based on extremely large cosmological volumes, \\eg 1 Gpc side-length. Such large volumes cannot be simulated at the resolution reached by observations, so that re-simulation techniques have been developed (\\eg GRAFIC \\citealt{1995astro.ph..6070B}; ZIC \\citealt{1997MNRAS.286..865T}; \\citealt{2010MNRAS.403.1859J}). When such high resolution is reached, the physic of the baryonic component must be followed with particular care. The magnetic field amplification, in particular, depends on the small scale motions of the gas. Hence, as the resolution increases smaller scale motions are revealed, and the magnetic field amplification increases accordingly \\citep[see \\eg ][]{2008SSRv..134..311D}. \\\\ In this paper we present a set of galaxy clusters extracted by a low resolution DM simulation and re-simulated at high resolution (the softening length is $\\sim$5 kpc $h^{-1}$) within a cosmological framework in order to resolve scales comparable to those reached by observations. This work is focused on the 24 most massive galaxy clusters ($M_{200}>10^{15} h^{-1} Mpc$) of our sample. Simulations are performed for the first time relaxing the assumption of ideal MHD, and including a resistivity term in the induction equation ($\\eta_m$). Our sample of simulated galaxy cluster is publically available for further studies\\footnote{contact a.bonafede@jacobs-university.de or kdolag@mpa-garching.mpg.de}. In this paper we present the simulated cluster sample: the MHD implementation with some test problems (Section. \\ref{sec:nonideal}), the re-simulation technique used (Section. \\ref{sec:clusters} and more detailed in the appendix); the effect of different values for $\\eta_m$ are analyzed and discussed in Section \\ref{sec:etam}, where the main properties of the clusters are also presented.  Finally, discussion and conclusions are reported in Section \\ref{sec:disc}.\\\\ This is a first paper aimed at presenting the cluster sample, the zoom-in initial conditions, and the non-ideal MHD implementation in the {\\tt GADGET} code. This sample has also been used by \\citet{2011arXiv1102.2903F} for a study of the scaling relations of X--ray mass proxies. In a future paper the authors will investigate in more detail the cluster properties, and the interplay between thermal and non-thermal components in the ICM. ",
        "conclusions": "\\label{sec:disc} We have presented a set of simulated galaxy clusters. It consists of 24 massive objects ($M_\\mathrm{vir}> 10^{15} h^{-1}$\\msun) re-simulated at high resolution up to 5-6 virial radii, plus 50 more clusters with $M_\\mathrm{vir}> 10^{14} h^{-1}$\\msun  that fell within this high resolution region.  This large set permits to study the cluster properties in a wide range of masses and at high resolution \\citep[see \\eg][]{2011arXiv1102.2903F}. The evolution of the clusters has been followed using the MHD implementation within the {\\tt GADGET-3} code \\citep{2009MNRAS.398.1678D}, that has been here modified in order to include the magnetic resistivity term in the induction equations. It is the first time that this term is analyzed in the context of cluster formation and evolution.  In this first paper we have presented the zoomed initial conditions, the non-ideal MHD implementation, and the average properties of the more massive clusters when a magnetic resistivity term is included in the MHD equations.  Further analysis will be performed in a future paper (Paper II, Bonafede et al. in prep) where the physical implications will be discussed in more detail.\\\\ Our main results can be summarized as follows: \\begin{itemize} \\item{Non-ideal MHD equations have been implemented within the {\\tt GADGET} code. The tests performed on two different problems show that the numerical implementation is accurate and can be used to study the effect of the magnetic resitivity.} \\item{The magnetic field profiles obtained with non-ideal MHD can reproduce the profile inferred from Faraday Rotation Measures of the Coma cluster. Four clusters having X-ray morphologies similar to the one of the Coma cluster have been selected to test the effect of changing the constant $\\eta_m$ used in the induction equation. The best agreement with the limits given by observations is achieved with $\\eta_m= 6\\times 10^{27} cm^2 s^{-1}$.} \\item{ The whole sample has been simulated using $\\eta_m=6\\times 10^{27} cm^2 s^{-1}$, and the derived magnetic field profiles are consistent with the Coma profile. The best-fit found for the Coma profile lies in fact between the rms of the simulated profiles.} \\item{We have fitted the magnetic field profile with a $\\beta$-like model, finding that the magnetic field profile of the simulated galaxy clusters can be well reproduced by values $B_0= 4.7 \\pm 1.7 $, and $\\mu = 0.46 \\pm 0.11$ (see Eq. \\ref{Eq:betamodel}), in good agreement with the value found for the Coma cluster. The value of $\\mu$ would correspond to a value of $\\alpha \\sim 0.6$ (Eq. \\ref{eq:Brho}) for a Coma-like gas density profile. } \\item{We have investigated possible correlations of the magnetic field strength with the cluster mass. The magnetic field strength, averaged over a central spherical volume of 0.3$R_\\mathrm{vir} h^{-1}$ in radius, is similar for all the clusters in the sample, in agreement with what has been recently found by \\citet{2011A&A...530A..24B}. This indicates that the presence of radio halo emission, found in a fraction of massive galaxy clusters, cannot be attributed to a difference in the magnetic field strength. A mild dependence of the magnetic field strength with cluster temperature is indicated by these simulations. } \\item{The density, temperature and entropy profiles of the simulated galaxy clusters have been derived for different values of $\\eta_m$. We find that the effect of a magnetic diffusive constant is visible in such profiles, leading to flatter temperature and entropy profiles in the inner region of the cluster ($R \\leq 0.1 R_\\mathrm{vir}$ at maximum). } \\end{itemize} The cluster sample and the new MHD-implementation we have presented is suitable to investigate other issues that are not discussed here, and that will be studied in a future paper, such as the interplay of the magnetic field with the thermal gas of the ICM (\\eg how is the thermal conduction modified, the role of the magnetic pressure in suppressing the cooling in the inner regions). In the next years, the LOw Frequency ARray (LOFAR) and the Expanded Very Large Array (EVLA) will allow us to improve our knowledge of the non-thermal component of the ICM, and a larger sample of data will be soon available for a more complete comparison."
    },
    "1107/1107.5583_arXiv.txt": {
        "abstract": " ",
        "introduction": "Neutrino telescopes, such as IceCube~\\cite{Ahrens:2003ix} or ANTARES~\\cite{Aslanides:1999vq}, are expected to reveal the nature of the sources of the highest-energetic cosmic rays. If a substantial amount of protons (or heavier nuclei) are accelerated in the sources, charged pion production is expected via $pp$ and $p\\gamma$ interactions, and therefore a neutrino flux.  There are numerous candidate sources, see \\Ref~\\cite{Becker:2007sv} for a review and \\Ref~\\cite{Rachen:1998fd} for the general theory. We focus on the prompt emission of gamma-ray bursts (GRBs) in this study, where photohadronic ($p\\gamma$) interactions are expected to lead to a significant flux of neutrinos~\\cite{Waxman:1997ti}.  If GRBs are the sources of the highest-energetic cosmic rays, the expected neutrino flux is just around the corner. In fact, recent IceCube-40 data (referring to the 40 string configuration of IceCube) start to test this paradigm~\\cite{Abbasi:2011qc,Ahlers:2011jj}, based on analytical estimates of the neutrino fluxes, see \\Refs~\\cite{Waxman:1997ti,Guetta:2003wi,Becker:2005ej,Razzaque:2006qa} and \\Ref~\\cite{Abbasi:2009ig} for the application to IceCube data. On the other hand, numerical calculations indicate that additional processes in the photohadronic interactions, flavor mixing, and magnetic field effects affect normalization and shape of the expected neutrino fluxes, see \\Refs~\\cite{Murase:2005hy,Lipari:2007su,Baerwald:2010fk}. In \\Ref~\\cite{Baerwald:2010fk}, we have made the impact of these effects on the original Waxman-Bahcall (WB) flux shape~\\cite{Waxman:1997ti} very explicit. For example, a characteristic multi-peak structure from high-energy processes in the photohadronic interactions, magnetic field effects, and flavor mixing has been predicted for the same astrophysical assumptions as for the WB flux.  While \\Ref~\\cite{Baerwald:2010fk} is essentially based on one set of parameters only, we focus in this study on {\\em aggregated} neutrino fluxes. Because of the low statistics expected from a single burst, any state-of-the-art analysis is based on such an aggregated flux, which can either be a diffuse flux, for which individual sources are not resolved, or a stacked flux, for which the neutrino fluxes from individual sources are expected to be correlated with their gamma-ray counterparts. The interpretation of aggregated flux limits, of course, implies new systematics not present for a single burst, such as by the integration over parameter distributions (for diffuse fluxes), or the assumptions for unknown burst parameters (stacked fluxes). In the stacking case, the time- and directional-wise correlation can be used to suppress backgrounds, and the fluence of the gamma-rays can even be used to compute the expected shape of the neutrino spectrum~\\cite{Abbasi:2011qc}. However,  since for many bursts the parameters needed for such an estimate are unknown, such as the redshift or Lorentz factor, ``standard values'' are often used. It is one of the motivations for this study to critically review these standard values from a theoretical perspective, in order to reduce the systematical errors in the interpretation of future analyses. For instance, depending on the method, a too large standard value for the redshift may easily lead to an overestimated neutrino flux.  In addition, we discuss if the spectral multi-peak features are present in aggregated fluxes. Note that experiments as IceCube have limits with optimal sensitivities in specific energy ranges, see \\eg\\ discussion in \\Refs~\\cite{Halzen:2010yj,Winter:2011jr}. Therefore, it is relevant if these spectral features are present in aggregated fluxes, especially in a diffuse flux.  Predictions for diffuse neutrino fluxes have, for instance, been made in \\Refs~\\cite{Halzen:1999xc,Murase:2005hy,Gupta:2006jm}. For a diffuse flux, typically at least the redshift distribution is integrated over, which is assumed to follow the star formation rate in some form. Apart from the redshift, interesting parameters which vary from burst to burst are Lorentz factor, luminosity, gamma-ray fluence, and variability timescale. For example, in \\Ref~\\cite{Gupta:2006jm}, relying on conventional fireball phenomenology, a log normal distribution for the Lorentz factor is assumed. In the conventional fireball model, the main contribution to the neutrino flux is assumed to come from small Lorentz factors, see \\Refs~\\cite{Guetta:2000ye,Guetta:2001cd} for discussions on the parameter space constraints. This conclusion is based on the estimate of the size of the interaction region from the variability timescale, which strongly depends on the Lorentz factor. However, in a modified version of the fireball model including with the hypothesis that the bursts have similar properties in the comoving frame (shock rest frame) the dominance of the small Lorentz factors is no longer given, as we will demonstrate. In fact, a recent analysis by Ghirlanda et al.~\\cite{Ghirlanda:2011bn}, which appeared during completion of this work, favors this hypothesis from the simple argument that the variations of the bulk Lorentz factor imply the Amati~\\cite{Amati:2002ny} and Yonetoku~\\cite{Yonetoku} relations, while the bursts are expected to be similar in the comoving frame. In order to clarify the model-dependence  (\\cf, \\Sec~\\ref{sec:source}), we start with the minimal ingredients (and input parameters) to neutrino production needed to numerically reproduce the analytical estimates in  \\Refs~\\cite{Waxman:1997ti,Guetta:2003wi,Abbasi:2011qc}. We clearly separate these from model-dependent assumptions (\\Sec~\\ref{sec:obs}), which has the advantage that one can clearly see where the model-dependent hypotheses enter. The main questions which we pursue in this context are: Assuming theoretical expectations for parameter distributions, from which parameters do the main contributions to the neutrino flux come from? As a consequence, what value should one assign to a burst with unknown redshift or Lorentz factor, such as in a stacking analysis?  What are the assumptions going into that conclusion, in particular, which assumptions hold in a model-independent way? In order to discuss these questions in a coherent and transparent way, we discuss each of the relevant input parameters separately in \\Sec~\\ref{sec:diffuse}.  In principle, a neutrino telescope cannot only measure the neutrino flux, but also distinguish different flavors by different event topologies~\\cite{Beacom:2003nh}. In the simplest case, one may use two topologies: tracks (mainly from $\\nu_\\mu$) and cascades (mainly induced by $\\nu_e$ and $\\nu_\\tau$) to construct an observable muon track to cascade flavor ratio as a flavor dependent quantity. Note that this flavor ratio is defined in the spirit of cascade searches from extragalactic neutrino sources, which has been very recently conducted in \\Ref~\\cite{IcCascade:2011ui} for IceCube-22. From the charged pion $\\pi \\to \\mu \\to e$ decay chain, the flavor composition at the source is given by  $\\nu_e:\\nu_\\mu:\\nu_\\tau = 1:2:0$ (``pion beam source''), and, for GRBs, changes to $0:1:0$ at higher energies~\\cite{Kashti:2005qa,Lipari:2007su} (``muon damped source'') because of the synchrotron cooling and decay of the intermediate muons, pions, and kaons. Although one may not expect to observe this effect with high statistics on the timescale of ten years of full IceCube running due to the lower effective area for cascades, some conclusions on the astrophysical source (such as the magnetic field~\\cite{Hummer:2010ai}) or the particle physics during neutrino propagation (see, \\eg, \\Ref~\\cite{Pakvasa:2008nx} for a short review) may be obtained at IceCube or future potential upgrades. Since the flux normalization drops out of this flavor ratio, it ought to be relatively robust with respect to astrophysical uncertainties~\\cite{Mehta:2011qb}.  In this study, we do not only discuss the muon neutrino flux at the detector, but also the flavor ratio. In particular, we will demonstrate that the aggregated neutrino flavor ratio, though perhaps measurable with much lower statistics, is a more reliable prediction for GRBs than the flux itself; see \\Sec~\\ref{sec:diffuse}.  Finally, we discuss the systematical error on the quasi-diffuse flux introduced by the stacking statistics, \\ie, the number of stacked bursts, in \\Sec~\\ref{sec:lowstatistics}. Note that this systematical error is important for the interpretation with respect to the paradigm that GRBs are the sources of the highest energetic cosmic rays. It is also illuminating to show what the main qualitative differences between stacked and diffuse fluxes are, and where we expect the borderline. Throughout this paper, we use a Monte Carlo method to compute the neutrino fluxes. Instead of  Monte Carlo integration over the parameter distributions, we compute the neutrino spectrum burst-by-burst, sampling over the input parameter distributions. This method has the advantage that it converges to conventional (Monte Carlo) integration in the limit of a large number of bursts, such as $\\mathcal{O}(10 \\, 000)$ bursts expected in the observable universe over a timescale of ten years. As we shall see, this limit then also represents a diffuse flux. On the other hand, small sample sizes can be chosen to simulate the impact of stacking analyses of \\eg\\ $\\mathcal{O}(100)$ bursts, as used in \\Ref~\\cite{Abbasi:2011qc}. ",
        "conclusions": "We have numerically computed GRB neutrino fluxes and flavor ratios including different meson production modes, magnetic field effects, and flavor mixing. Our starting point has been the minimal set of assumptions needed to compute the neutrino fluxes at the same level as it is used in recent IceCube stacked limits, but including the full spectral dependencies and the cooling of the secondaries explicitely. We have kept these assumptions as model-independent as possible. As the next step, we have connected the quantities needed for the neutrino production to different models, such as a cannonball-like model and conventional fireball phenomenology, where we have clearly demonstrated where model-dependent assumptions or relationships enter.  The main purpose of this study has been the computation of aggregated fluxes, where we have distinguished diffuse fluxes and quasi-diffuse fluxes obtained from stacking analyses. For the computation of the aggregated fluxes, we have computed the flux burst by burst, where we have used $10 \\, 000$ bursts for the diffuse flux -- as may be expected in the observable universe within ten years of neutrino telescope operation. We have established that these $10 \\, 000$ bursts lead to a flux close enough to the diffuse limit, \\ie, that the systematical error coming from the burst statistics is small in that case. Quantitatively, from the redshift distribution (strong evolution case), the neutrino flux in these ten years will be within 20\\% of the diffuse limit with 97\\% probability.  Our computations of aggregated fluxes have been based on parameter distributions for redshift $z$, magnetic field $B'$, Lorentz boost $\\Gamma$, and the spectral shape of the target photons. We have identified two major issues: The theoretical distribution of the parameter, and the interpretation of a different parameter value in terms of the relative contribution to the flux. A relatively simple case is redshift, which may follow the star formation rate in some form, and the interpretation of the relative contributions of the bursts, if assumed to be otherwise comparable at the source, is straightforward. A completely different case is $\\Gamma$ for which the theoretical distribution is unknown and the translation of the relative contribution of the bursts is model-dependent. For example, in the conventional fireball model, the  size of the interaction region is estimated from the observed variability timescale $t_v$, which implies that it strongly depends on the Lorentz factor $\\Gamma$. As a consequence, small values of $\\Gamma$ are assumed to dominate the neutrino flux. In the same model, the hypothesis of similar properties in the comoving frame leads to the dominance of high $\\Gamma$, as it is consistent with the recent Ghirlanda et al. study~\\cite{Ghirlanda:2011bn}. As a consequence, some of the \\textit{Fermi}-LAT accessible bursts with high $\\Gamma$ may contribute stronger to the neutrino flux than previously anticipated. Since the underlying assumptions for the different parameter distributions are qualitatively different, we have not mixed these different cases in the computation of the diffuse fluxes; in addition, our quantitative statements (see below) are based on the redshift distribution only, as the most conservative assumption.  The synchrotron cooling versus decay of the secondaries (pions, muons, kaons) in combination with multi-pion photohadronic processes leads to a characteristic multi-peak structure of the GRB neutrino flux. In addition, a transition  of the flavor composition from a pion beam source (neutrinos from pion decay chain) to a muon damped source (neutrinos from muon decays suppressed) at high energies is expected, which can, in principle, be measured by the muon track to cascade ratio in a neutrino telescope. We have established that these features are, in general, also present in diffuse fluxes, such as if the redshift of the bursts is varied. Only too strong variations of the magnetic field may destroy the spectral peaks in the flux and smear out the flavor ratio transition. In this case, the spectral shape of the flux cannot be used anymore to obtain information on the magnetic field, whereas the flavor composition depends on the mean value of $B'$ and the broadness of its distribution with no other significant parameter dependencies. This means that the flavor ratio can be regarded as the most robust prediction one can make for diffuse GRB neutrino fluxes. The observation of magnetic field effects can also be useful as a model discriminator: we have demonstrated that the observation of spectral features in diffuse GRB neutrino fluxes or a narrow transition of the flavor ratio can discriminate between the conventional fireball assumption that bursts with low $\\Gamma$ dominate the neutrino flux and the hypothesis that the bursts are alike in the comoving frame.  Apart from diffuse fluxes, we have also discussed quasi-diffuse limits obtained from the stacking of a relatively  small number of bursts, such as $\\mathcal{O}(100)$ bursts used in recent IceCube analyses. The stacking statistics of a small number of bursts has been shown to introduce a systematical error in the computation of the quasi-diffuse flux, which is at the level of about 50\\% for 100 bursts, 35\\% for 300 bursts, and 25\\% for 1000 bursts (90\\% CL) from the redshift distribution only -- additional parameter variations will lead to a larger effect (see \\App~\\ref{sec:threshold}). Therefore, $\\mathcal{O}(1000)$ bursts may have to be stacked for reliable conclusions. The reason is that the redshift $z$ leading to the largest contribution to the flux does not coincide with the maximum of the redshift distribution. As another consequence, in stacking analyses, it may also be consistent to assign $z \\simeq 1$ to a redshift not measured for a burst, since, depending on the method, a different value may lead to an overestimate of the neutrino flux or to wrong energies of the spectral neutrino breaks. This conclusion holds almost independently of the redshift distribution, including the strong evolution case. Similar standard values can be inferred for $\\Gamma$, which are, however, model dependent. For example, $\\Gamma \\simeq 700$ gives the main contribution in our cannonball ansatz, $\\Gamma \\simeq 400$ in the fireball model with the parameters fixed in the comoving frame, and  $\\Gamma \\simeq 140$ for the conventional fireball assumption. While normally a high stacking statistics is desirable, we have also demonstrated that indeed the fluctuations among different stacking samples may be used as a discriminator between cannonball and fireball approaches.  A very interesting special case is the observation of a single bright burst detectable without multi-messenger counterpart. From the redshift distribution (strong evolution case), we have estimated that the probability to observe at least one single bright burst individually detectable in IceCube during ten years of full-scale operation is about 40\\% if the Waxman-Bahcall flux is saturated, and 2\\% if the diffuse neutrino flux is an order of magnitude smaller. Thus, at best, there may be a 50-50 chance for such a single burst detection.  We conclude that neutrino observations of GRBs, both in terms of flux and flavor ratio, may not only help to test the paradigm if GRBs are the sources of the highest energetic cosmic rays, but also give important clues on the underlying models. Especially spectral features and the flavor composition turn out to be very interesting discriminators of very basic assumptions if a signal is observed. If no signal is observed, a systematical error coming from the stacking statistics has to be regarded in the context of the limit calculation.  \\subsubsection*"
    },
    "1107/1107.5060_arXiv.txt": {
        "abstract": "The epoch of galaxy assembly from $2\\leq z \\leq 4$ marks a critical stage during the evolution of today's galaxy population. During this period the star-formation activity in the Universe was at its peak level, and the structural patterns observed among galaxies in the local Universe were not yet in place. A variety of novel techniques have been employed over the past decade to assemble multiwavelength observations of galaxies during this important epoch. In this primarily observational review, I present a census of the methods used to find distant galaxies and the empirical constraints on their multiwavelength luminosities and colors. I then discuss what is known about the stellar content and past histories of star formation in high-redshift galaxies; their interstellar contents including dust, gas, and heavy elements; and their structural and dynamical properties. I conclude by considering some of the most pressing and open questions regarding the physics of high-redshift galaxies, which are to be addressed with future facilities. ",
        "introduction": "\\label{sec:intro}    Understanding the detailed formation and evolution of the galaxies we observe today remains one of the great challenges of modern cosmology. An exceedingly rich variety of galaxy properties exists in terms of luminosity, mass, color, structure, gas content, heavy-element enrichment, and environment, many of which, in turn, are strongly correlated with each other. As reviewed by \\citet{blanton2009}, based on the latest generation of wide-field surveys of the local Universe, astronomers have constructed an exquisitely detailed and statistically robust description of the galaxy population {\\it today}. These results are crucial in terms of providing a boundary condition, or endpoint, for our description of the formation and evolution of galaxies. Ultimately, we strive to tell that story from beginning to end. To do so, we must assemble additional observations and theories.   Multiple complementary approaches can be used to construct the history of galaxy formation. These include ab initio analytic models or numerical simulations; examinations of the fossil record contained in the ages, metallicities, and phase-space distributions of stars in nearby galaxies; and direct observations of distant galaxies, for which the cosmologically significant lookback time allows a probe of the Universe at an earlier time. In order to test theoretical models of galaxy formation at every time step, observations of galaxies over a wide range of lookback times are required. Furthermore, a robust translation must be performed between the observer's empirical quantities of luminosity, color, and velocity dispersion, and the theorist's physical quantities of stellar and dynamical mass, current star-formation rate and past star-formation history.   In the comparison between observations and theoretical models of galaxy evolution, an important recent development is the establishment of a precision cosmological framework. Observations of the cosmic microwave background radiation, large-scale structure, Type Ia supernova, the abundance of galaxy clusters, and the expansion rate of the Universe, all appear to be well described by a cosmological model in which the Universe is spatially flat, with the dominant component of the mass-energy density in the form of dark energy, and the remainder consisting mostly of cold dark matter with a small fraction of baryons. The initial spectrum of density fluctuations in this model is adiabatic, Gaussian, and nearly scale invariant. The most recent determination of cosmological parameters from the {\\it Wilkinson Microwave Anisotropy Probe} (WMAP) \\citep{spergel2003} is presented in \\citet{komatsu2011}, and highlights the fact that most parameters are determined with better than $5-10$\\% precision.  Constraining the background cosmological parameters is a crucial part of understanding galaxy formation not only for converting apparent quantities such as flux and angular size, respectively, into intrinsic ones such as luminosity, star-formation rate, stellar mass, and physical size. Cosmological parameters are also required for precise theoretical calculations whose predictions can be compared with observations, because, according to leading models, the underlying set of cosmological parameters determine how tiny primordial dark matter density fluctuations evolve under the influence of gravity into the large-scale spatial distribution of matter in the current Universe. In this framework, the collapsed perturbations of dark matter --  dark matter halos -- serve as the very sites of galaxy formation. Based on recent determinations of cosmological parameters, massive numerical simulations of the growth of dark matter structure have been performed \\citep{springel2005,boylankolchin2009}. In order to compare more directly with observations of galaxy formation, models (either numerical hydrodynamic or semi-analytic) describing the baryonic processes of gas cooling, star formation, and metal enrichment are also required, and these, too, are advancing with increased spatial resolution and complexity \\citep[e.g.,][]{ceverino2010,somerville2008}.   In order to probe the origin of the global patterns observed in the current galaxy population, we must look back to a time before these trends were already in place. Furthermore, the old stellar populations of nearby early-type galaxies in dense environments suggest that the bulk of their stars formed at $z\\geq 2$ \\citep{thomas2005}. Therefore, catching the formation of spheroids ``in the act\" requires observations at such early times. Given the strong correlation between the properties of the spheroidal components of galaxies and their central black holes, the epoch when the spheroids are forming holds special interest for explaining this connection. As described in more detail during the course of this review, the redshift range, $2\\leq z\\leq 4$, corresponding to a lookback time of $\\sim 10-12$~Gyr, is ideal for directly observing the progenitors of today's fairly luminous spheroidal and disk galaxies while in the very process of attaining the properties that come to define them over the next 10 Gyr up to the present day. Towards the end of this redshift range, the overall level of ``activity\" in the Universe -- both in terms of star formation and black hole accretion - was at its peak value. In contrast to what is observed in the present-day Universe, a significant fraction of the most massive galaxies still sustained active star formation. Furthermore, the Hubble Sequence of disk and elliptical galaxies was not yet in place, and the abundance of rich clusters was vanishingly small. The early Universe looked drastically different from its current state. Therefore, studying this epoch can yield important clues about the evolution of galaxies.   Within the past decade, there has been incredible progress in the study of galaxies in this important redshift range of galaxy assembly. A variety of novel techniques have been used to identify distant galaxies, and the sample of galaxies with spectroscopic redshifts at $2\\leq z\\leq 4$ now numbers  well into the thousands. Although we are far from approaching the overwhelming statistical power of the giant local redshift surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{abazajian2009,blanton2009} and 2dF Galaxy Redshift Survey \\citep[2dFGRS;][]{colless2001}, in terms of number of galaxies surveyed spectroscopically, volume probed, and data quality, key features of the galaxy population at these early times are emerging. The physical properties inferred for distant galaxies have provided important inputs and daunting challenges to state of the art theoretical models of galaxy formation. As we look forward to the next generation of instrumentation on current and future large ground-based telescopes, as well as the James Webb Space Telescope ({\\it JWST}) in space, and ever more sophisticated galaxy formation models, it is worth reviewing what is known about the physical properties of high-redshift galaxies at $2 \\leq z \\leq 4$.   This primarily observational review is constructed as follows. In Section~\\ref{sec:technique}, we provide an overview of the many different techniques that have recently been employed for identifying high-redshift galaxies. We continue in Section~\\ref{sec:empirical} by reviewing the global multiwavelength distributions in luminosity and color for galaxies in this redshift range. In Section~\\ref{sec:stellarpop}, we delve into the techniques used to transform empirical quantities such as luminosity and color into physical ones relating to galaxy stellar populations. In particular, we focus here on stellar content and the history of star-formation activity, as well as the relationship between these quantities.  Sections~\\ref{sec:ISM} and \\ref{sec:structure} in turn consider what is known about the interstellar medium (ISM) of distant galaxies --  where the interstellar contents include gas, dust, and metals -- and their structural properties and dynamics. Although many gaps remain in our knowledge of these fundamental physical properties, as described in Section~\\ref{sec:future}, future facilities and instrumentation will guide us in our quest to assemble a comprehensive picture of the galaxy population during this distant yet intriguing epoch in the history of galaxy formation. ",
        "conclusions": "\\label{sec:future}  Our path has traveled through many different wavelength ranges and observational techniques along the way to characterizing the stars, dust, gas, heavy elements, structure and dynamics of high-redshift galaxies at $2\\leq z \\leq 4$. By design, this review has focused primarily on the translation of observed quantities to physical ones, instead of making systematic comparisons with particular theories (although at times, when appropriate, a connection was made). At the same time, the rapid development of observations and their physical interpretation over the previous decade has yielded an important body of data for input into models of galaxy formation (i.e., the diversity of star-formation histories and structures among the highest-mass systems; the connection between stellar mass and star-formation rate; the nature of dust extinction as a function of luminosity;  the estimate of star-formation efficiency for LIRGs and ULIRGs; the mass-metallicity relation as a function of redshift; the compactness of massive, quiescent systems; the high degree of turbulence in the ISM of star-forming galaxies; and so on).  There are some definite limitations in the nature of this review. Indeed, the description of physical properties was often still couched in terms of how the results applied to a specific galaxy sample, assembled using a specific selection technique from among the ones described in Section~\\ref{sec:technique}. In order for measurements of galaxy properties to have true discriminating power among galaxy formation models, it is crucial to use samples that are complete with respect to a well-defined property, such as rest-frame optical or near-IR luminosity, stellar mass, dynamical mass, or star-formation rate. The majority of the samples described in this review slice the underlying high-redshift galaxy population in particular ways that add an extra layer of complexity if a comparison with a simulated galaxy population is desired. We advocate the design of future surveys with the type of physical completeness described above in mind. The NEWFIRM Medium-Band Survey \\citep{vandokkum2009_pasp} represents an important step in this direction, but an even deeper survey would be desirable, with both rest-UV and optical spectroscopic follow-up.   Furthermore, we point out that many of the results reported here were (by necessity) based on small samples -- these include mid-IR spectra of the most luminous sources, measurements of molecular gas content and dynamics, individual emission and absorption-line metallicity measurements at $z\\geq 2$, high-resolution rest-frame optical morphological measurements, and AO-assisted emission-line maps. While these results highlight truly compelling questions, they also await confirmation from samples an order of magnitude larger for a robust comparison with models.  Also, by necessity, some of the most intriguing results regarding gas, dust, and structural properties are limited to the luminous extreme of the galaxy population. With the steep luminosity functions described in Section~\\ref{subsec:empirical-LF}, galaxies with $L\\leq 10^{12} L_{\\odot}$ comprise the bulk ($\\sim 75$\\%) of the luminosity and star-formation rate density of the Universe at $z\\sim 2$. Even galaxies with $L\\leq 10^{11} L_{\\odot}$ may contribute $\\sim 30-35$\\% of the bolometric luminosity density at this redshift \\citep{reddy2008}. We must find ways of extending our physical studies towards fainter luminosities -- either by using gravitationally-lensed objects, or the next generation of instruments and telescopes. These faint objects are important, and perhaps more analogous to the galaxies playing a crucial role in the reionization process at $z\\geq 6$.  We close by highlighting two important observational challenges, and looking towards the future. First, the direct measurement of gas inflow (accretion) and outflow (star-formation and AGN feedback) received only passing reference during the course of this review. The ubiquity of galaxy-scale outflows among $z\\geq 2$ UV-selected galaxies has long been known on the basis of rest-frame UV and optical spectra \\citep{pettini2001,adelberger2003, shapley2003,steidel2004,adelberger2005,steidel2010}, yet obtaining robust constraints on the physical properties associated with these outflows (e.g., mass outflow rates) remains a challenge \\citep[but see, e.g.,][]{steidel2010}. At the same time, there is no obvious connection between these observations and the extremely popular theoretical models of cold gas accretion \\citep[e.g.,][]{keres2005,keres2009,dekel2009_nature}. An open challenge is to observe a ``smoking gun\" of gas accretion in high-redshift galaxies. Second, the study of galaxy environments at $z\\geq 2$ is a field in its infancy. While overdensities of LAEs have been identified near luminous radio galaxies \\citep{venemans2007}, and a small number of apparent protoclusters have been discovered serendipitously during the course of high-redshift galaxy spectroscopic surveys \\citep{steidel1998,steidel2005}, the study of the dependence of galaxy properties on environment at $z\\geq 2$ remains largely untapped.  Environmental studies will require extensive spectroscopy of mass-complete samples in both overdense and ``field\" environments, and will have potentially fundamental implications for understanding the origin of local galaxy environmental trends. Finally, we look forward to the truly exciting insights that will be made possible by upcoming facilities, including sensitive multi-object near-IR spectrographs on $8-10$-meter class telescopes, ALMA, JWST, and the extremely large ground-based telescopes of the next decade.   \\bigskip"
    },
    "1107/1107.0645_arXiv.txt": {
        "abstract": "T Tauri stars are young, low mass, pre-main sequence stars surrounded by an accretion disk. These objects present strong magnetic activity and powerful magnetic reconnection events. Strong shocks  are likely associated with fast reconnection in the stellar magnetosphere. Such shocks can accelerate particles up to relativistic energies. We aim at developing a simple model to calculate  the radiation produced by non-thermal relativistic particles in the environment of T Tauri stars. We want to establish whether this emission is detectable at high energies with the available or forthcoming $\\gamma$-ray telescopes. We assume that particles (protons and electrons) pre-accelerated in reconnection events are accelerated at shocks through Fermi mechanism and  we study the high-energy emission produced by the dominant radiative processes. We calculate the spectral energy distribution of T Tauri stars up to high-energies and we compare the integrated flux obtained with that from a specific {\\it Fermi} source, \\bicho, that we tentatively associate with this kind of young stellar objects (YSOs). We suggest that under reasonable general conditions nearby T Tauri stars might be detected at high energies and be responsible for some  unidentified {\\it Fermi} sources on the Galactic plane. ",
        "introduction": "\\label{s_intro} The \\textit{Fermi Gamma-Ray Space Telescope} was launched on 2008 June 11 and is the successor of the \\textit{Compton Gamma Ray Observatory}. It has two $\\gamma$-ray instruments on board, the Large Area Telescope (LAT) and the Gamma-ray Burst Monitor (GBM). The first one, a pair-production telescope, is the successor of the Energetic Gamma Ray Experiment Telescope (EGRET). LAT has a sensitivity $\\sim$ 10$^{-9}$ ph s$^{-1}$ cm$^{-2}$ in the energy range from 20 MeV to over 300 GeV.  Based on the first eleven months of survey, the {\\it Fermi} First Year Catalog has been published (Abdo et al. 2010). The catalog consists of 1451 sources  with  statistical significances of 4${\\sigma}$ or higher. This catalog provides potential identifications, when available. Nearly 630 {\\it Fermi} sources remained unidentified (Abdo et al. 2010). A fraction of these unidentified sources lay on the Galactic plane (latitudes $\\leq$ $5^{\\circ}$). They may be pulsars, supernova remnants (SNR), the effect of strong winds from massive stars, high-mass or low-mass X-ray binaries (HMXBs or LMXBs), or stellar clusters, among other possibilities (e.g. Romero, Benaglia, \\& Torres 1999).  Munar-Adrover, Paredes, \\& Romero (2011) have crossed the {\\it Fermi} First Year Catalog with catalogs of known YSOs, in order to identify those protostars that might emit $\\gamma$-rays. They conclude that 72\\% of the candidates obtained by spatial correlation should be $\\gamma$-ray sources with a confidence above 5$\\sigma$. Massive YSOs have already been  claimed to be $\\gamma$-ray  sources (e.g. Araudo et al. 2007, Bosch-Ramon et al 2010). However, such a claim has not been done yet for low-mass protostars.        In this paper we shall argue that these stars, in particular T Tauri stars, can be faint but sometimes detectable $\\gamma$-ray sources. We will focus our attention on the physical processes that can generate $\\gamma$-ray emission in T Tauri magnetospheres. Specifically, we shall discuss whether these protostars can be responsible for sources like 1FGL J1625.8$-$2429c. In the next section we present a brief introduction to T Tauri stars and we describe the general physical scenario. In Section 3 we provide the details of our  model.  In Section 4 we present our calculations  applied to a specific group of T Tauri stars and the main results. In Section 5,  we discuss these results and, finally, we close with some brief conclusions. ",
        "conclusions": "We have found that under some assumptions T Tauri stars might be  responsible for some nearby  {\\it Fermi} sources.  We have presented  a simplified model for the high-energy emission of this type of stars, that agrees with the available multiwavelength observations. T Tauri stars might be a new  class of galactic $\\gamma$-ray  sources in the Galatic plane. Based on this new scenario, 1FGL J1625.8-2429c is the first  candidate for collective $\\gamma$-ray emission from low-mass protostars. If the association with the $\\rho$ Ophiuchi cloud is confirmed, it would  be the closest $\\gamma$-ray source to the Solar System. This statement still holds even in the alternative case where the detected $\\gamma$-rays were simply due to cosmic rays interacting with the cloud ambient gas.    More complex  models can be developed in the future, using new and deeper observations. The energy where the cutoff lays might be established in the future by new  Cherenkov  instruments like CTA."
    },
    "1107/1107.5256_arXiv.txt": {
        "abstract": "We identify an abundant population of extreme emission line galaxies (EELGs) at redshift $z\\sim 1.7$~in the Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey (CANDELS) imaging from \\textit{Hubble Space Telescope}/\\textit{Wide Field Camera 3} (HST/WFC3).  69 EELG candidates are selected by the large contribution of exceptionally bright emission lines to their near-infrared broad-band magnitudes. Supported by spectroscopic confirmation of strong [OIII] emission lines -- with rest-frame equivalent widths $\\sim1000\\rm{\\AA}$ -- in the four candidates that have HST/WFC3 grism observations, we conclude that these objects are galaxies with $\\sim 10^8~\\msol$ in stellar mass, undergoing an enormous starburst phase with $M_*/\\dot{M}_*$ of only $\\sim 15$ Myr.  These bursts may cause outflows that are strong enough to produce cored dark matter profiles in low-mass galaxies.  The individual star formation rates and the co-moving number density ($3.7 \\times 10^{-4} ~ \\rm{Mpc}^{-3}$) can produce in $\\sim$4 Gyr much of the stellar mass density that is presently contained in $10^8-10^9~\\msol$ dwarf galaxies.  Therefore, our observations provide a strong indication that many or even most of the stars in present-day dwarf galaxies formed in strong, short-lived bursts, mostly at $z>1$. ",
        "introduction": "The formation history of dwarf galaxies with masses $\\sim 10^8~\\msol$ can usually only be studied through 'archaeological' age reconstruction, based on resolved stellar populations \\cite[e.g.,][]{grebel97, mateo98, weisz11}.  Their high-redshift progenitors have so far remained elusive despite the ever increasing depth of spectroscopic observing campaigns and imaging from the ground and the \\textit{Hubble Space Telescope} (HST).  In this paper we identify an abundant population of $z>1$ dwarf galaxies undergoing extreme starbursts, through HST/\\textit{Wide Field Camera 3} (WFC3) imaging from the Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey \\citep[CANDELS,][]{grogin11, koekemoer11}, that may well be the progenitors of present-day dwarf galaxies with stellar masses $\\sim 10^8-10^9~\\msol$.  At the present day, starbursts contribute a minority to the total star formation activity in dwarf galaxies \\citep{lee09}.  However, there is abundant evidence that the star formation histories are complex and that bursts play an important role \\citep[as reviewed by][]{mateo98}. Many authors find evidence for short-lived ($\\sim$10 Myr) SF events in nearby star-forming dwarf galaxies from a range of observational and modeling techniques \\citep[e.g.,][]{schaerer99, mas99, thornley00, tremonti01, harris04}, while others argue that star formation epochs of dwarf galaxies are more prolonged \\citep[e.g.,][]{calzetti97, lee08, mcquinn09}.  Simulations also indicate that star formation histories of low-mass galaxies are episodic or even burst-like \\citep[e.g.,][]{pelupessy04, stinson07, nagamine10}.  As most stars in dwarf galaxies formed more than 5 Gyr ago \\citep[e.g.,][]{dolphin05, weisz11}, it is crucial to understand the mode of star formation in dwarf galaxies at those early epochs, but 'archaeological' studies do not have the resolution in terms of stellar population age to constrain strengths, durations, and frequency of bursts.  The increased frequency of interactions with other galaxies and higher gas fractions at $z>1$ may have resulted in strong, short-lived starbursts.  In this paper we place the first constraints on the open question of how many and how frequently strong, short-lived starbursts occur in dwarf galaxies at $z>1$, and how relevant this mode of star formation is for the build-up of the dwarf galaxy population in a cosmological context.         \\begin{figure}[t] \\epsscale{1.2} \\plotone{f1.ps} \\caption{ Observed $I-J$ vs.~$J-H$ colors (in AB magnitudes) from HST/WFC3 and ACS imaging for all objects in the UDS and GSD (small black points) and the sample of emission-line dominated objects (large red points with error bars), selected by $I-J>0.44+\\sigma(I-J)$ and $J-H<-0.44-\\sigma(J-H)$, where $\\sigma(I-J)$ and $\\sigma(J-H)$ are the $1\\sigma$ uncertainties on the colors.  The blue line represents the redshifted ($z = 1.7$) continuum colors of the Starburst99 model \\citep{leitherer99} for continuous star formation, in the age range from 1 Myr to 100 Myr. The red line represents the same model, but with the $J$-band flux density increased by the emission line luminosity predicted by the model (Starburst99 predicts H$\\alpha$ luminosity -- [OIII] emission, which falls in the $J$ band at $z=1.7$, is assumed to have the same equivalent width). The black arrow indicates dust attenuation.} \\label{col} \\end{figure} ",
        "conclusions": "\\subsection{Comparison with Other Samples}  Galaxies with similar properties have previously been identified through broad-band photometry at $z<0.4$ in the Sloan Digital Sky Survey \\citep{cardamone09}, and have been shown by \\citet{amorin10} and \\citet {izotov11} to constitute the most strongly star-forming and most metal poor tail of the well-known class of blue compact dwarf galaxies \\citep[e.g.,][]{sargent70, thuan81, griffith11}, which have very low metallicities and extremely high, spatially concentrated star-formation activity \\citep{guzman98, overzier08}.  \\citet{cowie11} \\citep[also see][]{scarlata09}) recently studied the Ly$\\alpha$ properties of high-EW H$\\alpha$ emitters, providing a direct connection between higher-redshift searches of Ly$\\alpha$ \\citep[e.g.,][]{ouchi08, hu10}, and find Ly$\\alpha$ EWs ranging from 20$\\rm{\\AA}$ to 200$\\rm{\\AA}$.  Combining this with the findings of \\citet{nilsson11}, who show that Ly$\\alpha$ emitters at $z\\sim 2$ are objects with a very wide range in properties, it is clear that from Ly$\\alpha$ emitters one cannot derive a complete description of star formation in low-mass galaxies.  On the other hand, Ly$\\alpha$ emitters at higher redshifts ($z>3$) appear to be young, with small stellar masses \\citep[e.g.,][]{finkelstein09}, similar to the emission line galaxies studied here.  Narrow-band surveys identified galaxies with strong [OIII] and H$\\alpha$ emission lines with $\\rm{EW}\\sim 100-1000\\rm{\\AA}$ at redshifts $z=0.3-1$ \\citep[e.g.,][]{kakazu07}, demonstrated to be young and metal poor \\citep{hu09}.  Most notably, \\citet{atek10} pointed out the existence of a class of emission line galaxies at $z\\sim 1.5$ with $\\rm{EW}> 1000\\rm{\\AA}$ that would most likely be included in our sample as well.  However, so far, their nature has not been described and their cosmological relevance in the context of galaxy formation has remained unclear.  Therefore, let us now put these starbursting dwarf galaxies in a cosmological context.  \\subsection{Cosmological Context: Implications for the Formation of Dwarf Galaxies}  Our sample with redshifts $1.6 < z < 1.8$ consists of 69 low-mass ($\\sim 10^8~\\msol$), young ($\\sim 0.5-4 \\times 10^7~\\rm{yr}$), extreme starbursting, presumably metal-poor galaxies.  Their co-moving number density\\footnote{the values for our two widely separate fields, UDS and GSD, differ by only 12\\%} is $3.7\\times10^{-4}~\\rm{Mpc}^{-3}$, two orders of magnitude higher than that of nearby galaxies with similar EWs \\citep{cardamone09}.  The individual star formation rates and the number density combine into $1.7\\times10^{-3} ~\\msol ~\\rm{yr}^{-1} ~\\rm{Mpc}^{-3}$.  This is a 2\\% contribution to the total star-formation rate density at $z\\sim 1.7$ occurring in galaxies that contribute perhaps $\\sim 0.1\\%$ to the total stellar mass density at that epoch \\citep[e.g.,][]{karim11}.  Placing these detections in the context of the burst fraction among equally massive dwarf galaxies is difficult, however, since we cannot currently constrain their number density; galaxies without starbursts have older luminosity weighted ages and, as a consequence, are at least 2 magnitudes fainter than the starbursting galaxies.  With marginal detections, even in the latest WFC3 data, and no obvious spectral features, their redshifts cannot easily be estimated.  Model predictions differ strongly from observational measurements of the galaxy stellar mass function \\citep{guo11} and cannot be used to constrain the burst fraction among the population of low-mass galaxies.  Nevertheless, we can still gauge the importance of the observed burst for the formation of low-mass galaxies by making the reasonable (and testable) assumption that the observed bursts occur with the same frequency at all epochs $1\\lesssim z \\lesssim 3.5$ -- a period of $\\sim$4 Gyr during which the cosmic star formation history peaked and after which the number density of starbursts declines as mentioned above.  The basic indication that such bursts are important is that the number of stars produced in such bursts over a period of several Gyr is comparable to the number stars in present-day dwarf galaxies. Let us use this consideration to construct a simple toy model that relates the observations presented here to the mass function of present-day low-mass galaxies.  \\citet{guo11} use the data from \\citet{baldry08} for the galaxies with masses down to $10^7~\\msol$, and their mass function can be represented by a simple power law for all galaxies with masses $<10^{10}~\\msol$, that is, well below the knee of the \\citet{schechter76} function: \\begin{equation} \\phi(M_*) \\sim 0.043 \\Big(\\frac{M_*}{10^8~\\msol}\\Big)^{-0.37}, \\end{equation} in units of $\\rm{Mpc}^{-3}~\\rm{d}\\log(M_*)^{-1}$.  Let us then express the stellar mass of the present-day descendants of the observed starbursting galaxies at $z\\sim 1.7$ in terms of the following star formation history: \\begin{equation} M_{\\rm{desc}} = \\frac{M_{\\rm{burst}} \\times N_{\\rm{burst}}}{f_{\\rm{burst}}}, \\end{equation} where $M_{\\rm{burst}}\\sim 2 \\times 15~\\rm{Myr} \\times 5~\\msol~\\rm{yr}^{-1} = 1.5\\times 10^8~\\msol$ is the total stellar mass produced in a single starburst (the factor 2 is included to convert the observed burst age to its total duration), $N_{\\rm{burst}}$ is the number of bursts that occurs in each galaxy over the $\\sim$4 Gyr period between $z=1$ and $z=3.5$, and $f_{\\rm{burst}}$ is the fraction of the total stellar mass that is produced in such bursts -- the rest is assumed to form in smaller bursts and/or a more quiescent mode of star formation.  Now we can predict the number of present-day descendants by dividing the number of observed bursts at $z\\sim 1.7$ by the duty cycle, that is, the fraction of the time that a bursts is occurring over the period of time that they \\textit{can} occur (here assumed to be $1\\lesssim z \\lesssim 3.5$, or 4 Gyr).  This duty cycle is the ratio of the duration of a single burst ($2\\times 15$~Myr) and this 4 Gyr period, multiplied by the number of bursts per galaxy $N_{\\rm{burst}}$. Thus we have  \\begin{equation} \\phi(M_{\\rm{desc}}) \\sim \\frac{\\phi_{\\rm{burst}}}{N_{\\rm{burst}}} \\times \\frac{4~\\rm{Gyr}}{30~\\rm{Myr}}, \\end{equation} where $\\phi_{\\rm{burst}}$ is the co-moving number density of the bursts observed at $z\\sim 1.7$ ($3.7\\times10^{-4}~\\rm{Mpc}^{-3}~\\rm{d}\\log(M_*)^{-1}$, where we use the observed span in stellar masses of an order of magnitude -- see Figure 6 -- to introduce this differential unit).  According to Eqs.~2 and 4 the mass of the descendant, $M_{\\rm{desc}}$, is uniquely determined by the fraction of stars formed in bursts, $f_{\\rm{burst}}$.  Equating $\\phi(M_{\\rm{desc}})$ (Eq.~4) and $\\phi(M_*)$ (Eq.~2), and substituting for $N_{\\rm{burst}}$ (Eq.~3) we obtain, after rearranging terms: \\begin{equation} \\frac{M_{\\rm{desc}}}{10^8~\\msol} \\sim 2.4f_{\\rm{burst}}^{-1.6}. \\end{equation}  Thus, if $f_{\\rm{burst}}$ is close to unity, that is, almost all stars are formed in bursts, then we infer that each galaxy must undergo one or two bursts on average and that, therefore, $M_{\\rm{desc}} \\sim 1-2 M_{\\rm{burst}}$.  It is perhaps more realistic to adopt a smaller value for $f_{\\rm{burst}}$.  If $f_{\\rm{burst}}\\sim 0.5$, we find that two or three bursts must occur in each galaxy, producing descendants with masses $\\sim 10^{9}~\\msol$.  The latter would imply a growth in stellar mass between $z\\sim 1.7$ and the present by at least a factor 3 given the mass constraints on the underlying populations of the observed starburst galaxies.  Galaxy formation models suggest that the typical growth is indeed a factor of 3 or 4, but we should bear in mind that these models do not reproduce the observed co-moving number density evolution of low-mass galaxies with redshift \\citep[e.g.,][]{guo11}.  Therefore, these predictions should be treated with care.  Choosing $f_{\\rm{burst}}$ very low ($\\lesssim 0.1$) implies a very large growth in mass, with high-mass descendants ($>10^{10}~\\msol$). Such growth by, say, more than an order of magnitude in mass seems unlikely in the context of current galaxy formation models. In particular, models are better observationally constrained for these higher masses and the prediction is that more massive galaxies ($10^{10-11}~\\msol$) grow in mass by a factor 3 or 4, not by two orders of magnitude, between $z\\sim 1.7$ and the present \\citep[e.g.,][]{keres09}.  Given these constraints, the general conclusion we can draw is that our observations suggest that many or most stars in present-day dwarf galaxies (with masses $\\lesssim 10^9~\\msol$) have formed in a small number of starbursts at $z>1$.  The quantitative interpretation we offer here relies on our estimate of the characteristic burst duration (30 Myr).  However, the uncertainty in this assumed burst duration does not affect our conclusion that intense starbursts at high redshift constitute an important phase in the mass buildup low-mass present-day galaxies.  If the bursts last significantly longer, which we deem unlikely, then the total mass formed in a burst will be larger than assumed above, while the frequency of the bursts will remain the same since our selection method is only sensitive to young bursts.  As a result, the fraction of stars formed in bursts, $f_{\\rm{burst}}$ will be larger than what we derived above.  Significantly shorter bursts (say, $\\sim 5-10$~Myr, as allowed by the SB99 model -- see Sec.~4.2) are also unlikely, as the bursts are spatially extended, and the implied crossing time in these systems is about 30 Myr.  Nonetheless, in the case that burst times are short, the burst frequency must be higher to account for their observed number (Eq.~4), but the total mass produced in each burst remains the same (see Sec.~4.2).  As a result, the number of bursts per galaxy ($N_{\\rm{burst}}$) and, therefore, the fraction of stars formed in bursts ($f_{\\rm{burst}}$) increase.  The main assumptions in our interpretation, then, are 1) that the descendants of the observed galaxies remain low-mass galaxies up to the present day ($\\lesssim 10^9~\\msol$ in stars), and 2) that the observed bursts do not only occur at $z=1.6-1.8$, but are equally frequent over a much broader redshift range ($1\\lesssim z \\lesssim 3.5$). The first assumption is supported by the understanding that more massive galaxies are not expected to grow in stellar mass by several orders of magnitude, as would be required if the descendants of the observed starbursting galaxies are much more massive.  The second assumption is straightforward to test observationally by searches for similar objects over a wider redshift range.  \\subsection{Energy Budget and Core Formation}  The most remarkable property of these galaxies are their growth rates (specific star formation rates) of $20-200~\\rm{Gyr}^{-1}$.  This is far outside the realm of normal, more massive star forming galaxies, which typically have $\\sim 1~\\rm{Gyr}^{-1}$ \\citep[e.g.,][]{karim11}.  The amount of energy deposited into the interstellar medium through winds and supernovae ($10^{56-57}$~erg) exceeds the binding energy by an order of magnitude, even if the total mass is larger than the stellar mass by many factors.  This implies that the gas that is fueling the starburst may be in the process of being blown out, resulting in the end of the starburst phase.  Whether the gas will be blown out of the halo depends on the halo mass, which is currently unconstrained.  If the gas does not leave the halo it may eventually once more cool and sink to the center of the potential well and trigger another starburst.  Then, the observed bursts could be part of a semi-periodic cycle of star-formation activity, which is seen in simulations of low-mass galaxies \\citep[e.g.,][]{stinson07}.  Episodic star formation would also alleviate the difficulty of explaining the sudden occurence of such a large starburst in such small galaxies.  If such starbursts are not cyclical, then another possible explanation, usually offered to explain the lack of star formation at early times in even less massive systems, is that star formation had been suppressed at earlier epochs as a result of UV background radiation \\citep[e.g.,][]{babul92, babul96}.  An intriguing possibility is that if most of the gas is expelled from the central region, then the stars themselves could become unbound as well, dissolving the entire galaxy.  Displacement of large amounts of material leads to changes in the potential, and if this happens at very short time scales, then even the dark matter profile can be altered.  Outflows have long been argued to play a role in producing cored density profiles in low-mass galaxies \\citep[e.g.,][]{navarro96}.  The time scales, masses, and star-formation rates we derived above for the galaxies in our sample roughly meet the requirements for such a process to be efficient. Recently, \\citet{pontzen11} showed that episodic star formation, even with less intense bursts, can produce the same effect.  What the previous and subsequent evolution of the systems we observed is remains to be seen, but the energy balance and the possibility that such bursts are reoccurring make our observations consistent with this general picture; indeed outflows may be responsible for the formation of flattened dark matter profiles in low mass galaxies.  These speculative scenarios can be tested further with observational constraints on the gas masses and hydrodynamical modeling of these systems."
    },
    "1107/1107.0049_arXiv.txt": {
        "abstract": "In this paper, we address the controversy regarding the recent extended solar minimum as seen in helioseismic low- and intermediate-degree mode frequencies: studies from different instruments identify different epochs of seismic minima. Here we use mode frequencies from a network of six identical instruments,  Global Oscillation Network Group, continuously collecting data for more than 15 years, to investigate the epoch of minimum in solar oscillation frequencies  prior to the beginning of solar cycle 24. We include both low- and intermediate-degree modes in the  $\\ell$ range of 0 -- 120 and frequency range of 2.0 -- 3.5 mHz. In this analysis, we demonstrate that there were indeed two minima in oscillation frequencies, depending upon the degree of modes, or more precisely  the lower turning point radius of the propagating wave. We also analyze frequencies as a function of latitude to identify the beginning of solar cycle 24. We observe two minima at high latitudes and a single minimum at mid/low latitudes. This scenario is in contrast to cycle 23 where the epoch of seismic minimum did not change with latitude or depth. Our results also hint towards a possible role of the relic magnetic field  in modifying the oscillation frequencies of modes sampling deeper layers. ",
        "introduction": "The prolonged period of minimal solar activity and the delayed onset of solar cycle 24 have invoked a great deal of interest among solar physicists to identify precisely when the activity minimum occurred and how different it was from previously recorded minima. A variety of data from the solar interior to the corona has been analyzed to seek the origin of such an unusually low solar activity for a long duration (e.g., see articles in Cranmer, Hoeksema \\& Kohl 2010). All of these analyses  indicate that the current minimum was much deeper than previous ones in modern era. Since the source of solar activity is believed to lie in a shear layer at the base of the convection zone, known as the tachocline, the analysis of helioseismic data is important in order to probe these regions.  Fortunately, the availability of continuous, consistent helioseismic data for two consecutive solar minima has provided a unique opportunity to study the changes in the solar interior that might have led to this unusual minimum.  Using Birmingham Solar-Oscillation Network (BiSON) frequencies of low-angular degree ($\\ell \\le$ 3) modes for about three  solar cycles, \\citet{bison09} found that the frequencies were significantly lower in the extended minimum as compared to those during the previous ones. Similar results were obtained by \\citet{david09} in an analysis based on the data collected by the space-based  Global Oscillations at Low Frequencies (GOLF) on board the {\\it Solar and Heliospheric Observatory (SOHO)} spacecraft {\\citep{gab95}}. They further pointed out that cycle 24 started in late 2007, despite the absence of any visible activity on the solar surface. The analyses of modes in intermediate-degree range  between 20 and 150, obtained from the Global Oscillation Network Group (GONG: \\citet{harvey96}),  also showed lower oscillations frequencies during the recent minimum \\citep{jain10, sct10b}, however these studies did not identify the minimum as occurring until the end of 2008, leading to an extended phase of low activity. Furthermore, the zonal and meridional flow patterns inferred from inverting frequencies also suggested a delayed onset of the  new cycle. \\citet{howe09} compared the evolution of the zonal flow pattern in the upper convection zone and  suggested that the mid-latitude flow band corresponding to the recent cycle moved more slowly towards the equator than was observed in the previous cycle. As a result, there was a gradual increase in the apparent length of the cycle. \\citet{antia10} confirmed these results and compared the behavior of the low-latitude prograde band in past two minima. They pointed out that the flow band bifurcated in 2005 after merging in 2000, and merged again in 2007 before disappearing around the beginning of the new activity cycle in 2009. This is different from the previous minimum where split bands merged at the beginning of 1996 and disappeared soon after the activity minimum was reached. Evidence that the meridional circulation developed at medium-high latitudes before the beginning of a new cycle also supported the extended duration of low/no activity and delayed onset of cycle 24 \\citep{igh10}.  While most studies suggest that the beginning of cycle 24 occurred in late 2008/early 2009, the results from the low-degree mode frequencies were puzzling. In this context, we investigate the response of frequencies of low-degree as well as intermediate-degree modes during the recent extended period of low activity. The low-degree frequencies analyzed by \\citet{david09}  were obtained by unresolved (Sun-as-a-star) Doppler velocity observations, while intermediate-degree frequencies, studied by \\citet{sct10b}, were obtained with resolved  Doppler observations. To eliminate the influence, if any, of the techniques of observation and mode fitting, we use data from a network of six identical  instruments that treats both low- and intermediate-degree modes in a similar way. Finally, the depth and latitudinal dependence of the modes are studied during the last two minima. We present an analysis confined to modes that sense information from the tachocline region at different latitudes which provides useful information on the rise of a new cycle at mid latitudes. ",
        "conclusions": "It is been argued for a long time that the perturbations of near-surface layers generated by the changes at the tachocline are mainly responsible for the changes in frequencies, however the observations during the extended minimum provides an indication that there might be some effects from layers as deep as the core. The analysis of mode frequencies with lower turning points in the inner 30\\% of solar interior indicates a minimum earlier than that obtained in outer layers or surface activity. Further, \\citet{bison10} reported a quasi-biennial (2 year) signal in low-degree solar oscillation mode frequencies. Their study suggests that the 2 year signal, a predominantly additive contribution to the acoustic 11 year signal, has its origin in significantly deeper layers than the 11 year signal and is positioned below the upper turning point of the modes. Since the depth of a mode's upper turning point increases with decreasing frequency, one can expect low-$\\ell$ modes to be more influenced by the shallower layers. Thus, they argued that a second dynamo, seated near the bottom layer extending 5\\% below the solar surface, may be responsible for the 2 year signal.  Finally, the presence of two dynamos operating at different depths might be responsible for the observed double minima during the extended period of low activity.  To examine this quasi-biennial signal in the GONG frequencies, we follow the procedure as adopted by \\citet{bison10}; we subtract a smooth trend from mean shifts of independent time series by applying a boxcar filter of width 2 years. Figure~8 illustrates both the 11-year signal of the solar cycle and the frequency residuals for four different layers discussed in Figure~3. Although we find a weak 2-year periodicity at high activity period which is similar to the periodicity seen in proxies of solar activity \\citep{beno98}, no significant trend during the periods of low activity is noticed. We further apply the fast Fourier transformation to identify any significant trend in the frequency residuals. We did not find any prminent peak in this analysis which indicates that no other periodicity, except the 11-year solar cycle signal, is present in the GONG data. Thus, the absence of any significant trend of quasi-biennial signal (except during periods of high activity) suggests a different scenario that hints towards different physical processes which might be involved at different depths.  To investigate the solar origin of double minima seen in the oscillation frequencies, we consider two major fields inside the Sun: {\\it (i)} a weak field at the core that plays a crucial role in the generation of 22-year magnetic polarity cycle \\citep{mursula01} and is believed to have been present in the Sun since its formation \\citep{cowling45}, {\\it (ii)} a megagauss field just beneath the convection zone that is responsible for the dynamo mechanism \\citep{wad89} and the 11-year cyclic variation in solar activity. We speculate that a competition between these fields plays a significant role in explaining the variation in oscillation frequencies. It is possible that the relic field senses the minimum eariler than the field seated at the base of the convection zone. This is manifested in the form of early minimum seen in low-degree modes whose lower tuning points lie in the core. Further, the existence of two minima for the modes returning from radiative zone indicate an inter-connection between weak and strong magnetic fields. We anticipate that these fields affect the frequencies in all solar cycles but the influence due to relic field is shielded by the higher field strength. However, during the extended minimum phase, the magnetic field generated by the solar dynamo was relatively weak and hence the effect of the relic field could be observed.  The scenario presented here is speculative in nature since the information about the interior can only be obtained through indirect measurements, e.g. frequencies of solar oscillations. The continuing efforts to measure high-precision oscillation data for a complete 22-year magnetic cycle may unveil the influence of a relic field on the variation of oscillation frequencies."
    },
    "1107/1107.5941_arXiv.txt": {
        "abstract": "We report a first case study of the phase diagram of 2+1 flavor strongly interacting matter in $\\beta-$equilibrium, using the Polyakov$-$Nambu$-$Jona-Lasinio model. Physical characteristics of relevant thermodynamic observables have been discussed. A comparative analysis with the corresponding observables in the Nambu-Jona-Lasinio model is presented. We find distinct differences between the models in terms of a number of thermodynamic quantities like the speed of sound, specific heat, various number densities as well as entropy. The present study is expected to give us a better insight into the role that the superdense matter created in heavy ion collision experiments play in our understanding of the properties of matter inside the core of supermassive stars in the Universe. ",
        "introduction": "\\label{sc.intro}  The phase diagram of strongly interacting matter has been at the center of attention for quite some time now. Under a variety of extreme conditions of temperature and/or density the hadrons may overlap and loose their individuality and a new state of matter called Quark Gluon Plasma (QGP) may be formed \\cite{muller}. It is well believed that such a state of matter existed in the hot early Universe, a few microsecond after the Big Bang. Deconfined quark matter could also exist in the core of neutron stars (NS) \\cite{rajagopal-condmatt, blaschke-2010, glendenning3} where the temperature is relatively low but density is high. So an understanding of the physics of strongly interacting matter at such environmental conditions would have important cosmological and astrophysical significance.  In the laboratory such conditions of large temperatures and densities can be created by the collision of heavy ions at high energies. Presently the strongly interacting matter at high temperature and close to zero baryon densities $-$ a scenario relevant for early universe $-$ is being explored at Relativistic Heavy Ion Collider (RHIC) at BNL and the Large Hadron Collider (LHC) at CERN. A wealth of information has been obtained from RHIC, and a lot more is expected from both the future runs there as well as from LHC. More recently a variety of energy scans at RHIC and the upcoming facility (FAIR) at GSI, are expected to give us a glimpse of matter in the baryon-rich environments - the so-called {\\it compressed baryonic matter} (CBM). These experiments will also be useful in the search for signatures of critical phenomena associated with a second order critical end point (CEP).  At the same time, observational data are being collected by a large number of telescopes and satellites \\cite{bielich} such as the radio telescopes at the Arecibo, Parkes, Jo-drell Bank, and Green Bank Observatories, the Hubble Space Telescope, European Space Agency's International Gamma Ray Astrophysics Laboratory (INTEGRAL) satellite, Very Large Telescope (VLT) of the European Southern Observatory, the X-ray satellites Chandra, XMM-Newton and NASA's Rossi X-ray Timing Explorer and the Swift satellite. Observations from these facilities are supposed to tell us about the properties of strongly interacting matter at high densities relevant for the astrophysics of compact stars.  Thus on one hand the laboratory experiments are expected to scan the phase space temperature and various conserved quantum number densities of strongly interacting matter. On the other hand the astrophysical observations are expected to uncover the physics for high baryon number density region of the phase diagram. It should be noted here that the physical characteristics of the matter under consideration may be quite different in the two cases. The time-scale of the dynamics of heavy-ion experiments is so small that only strong interactions may equilibrate thermodynamically. While the dynamics in the astrophysical scenario is slow enough to allow even weak interactions may equilibrate. Thus a question naturally arises $-$ to what extent can laboratory experiments be used to infer about the compact star interiors? The aim of this paper is to address this question at a preliminary level from the characteristics of the \"$\\beta-$equilibrated\" phase diagrams.  One should be able to study the properties of systems described above from first principles using Quantum Chromodynamics (QCD), which is {\\it the} theory of strong interactions. However, QCD is highly non-perturbative in the region of temperature and density that we are interested in. The most reliable way to analyze the physics in this region of interest is to perform a numerical computation of the lattice version of QCD (Lattice QCD). The scheme is robust but numerically costly. Moreover, there are problems in applying this scheme for the systems having finite baryon density. Thus it has become a common practice to study the physics of strongly interacting matter under the given conditions using various QCD inspired effective models.  Until now various quark models, such as, different versions of the MIT bag model \\cite{burgio,alford1}, the color-dielectric model \\cite{ghosh_star1,ghosh_star2} and different formulations of the NJL model \\cite{buballa,baldo1} have been used to study the NS structure. Despite the similarity of the results on the value of the maximum NS mass, the predictions on the NS configurations can differ substantially from model to model. The most striking difference is in the quark matter content of the NS, which can be extremely large in the case of EOS related to the MIT bag model or the color-dielectric model, but it is vanishingly small in the case of the original version of the NJL model \\cite{buballa,schertler}. In the case of NJL model it turns out that, as soon as quark matter appears at increasing NS mass, the star becomes unstable, with only the possibility of a small central region with a mixed phase of nucleonic and quark matter. This may be a result of the lack of confinement in NJL model. In fact an indirect relationship between confinement and NS stability has been found in a study using NJL model with density dependent cut-off \\cite{baldo2}. Hence it is important to study the EOS from the Polyakov$-$Nambu$-$Jona-Lasinio (PNJL) model \\cite{Ratti1,ghosh-prd73,Ghosh:2008}, where a better description of confinement has been incorporated through Polyakov loop mechanism. Moreover, a comparison with NJL model might be helpful in understanding the role of Polyakov loop at high chemical potential.  A detailed study of 2+1 flavor strong interactions have been done by some of us using the PNJL model. The general thermodynamic properties along with the phase diagram \\cite{paramita}, as well as details of fluctuation and correlations of various conserved charges \\cite{anirban} have been reported. Here we extend the work by including $\\beta-$equilibrium into the picture. In the context of NJL model such a study was done earlier in \\cite{buballa-prd,hanauske-njl}. The properties of pseudoscalar and neutral mesons have been studied in finite density region within the framework of 2+1 flavor NJL model in $\\beta-$equilibrium\\cite{costa_meson1,costa_meson2}.  We investigate and compare different properties of the NJL and PNJL models in the T-$\\mu_B$ plane. The specialization of these studies to the possible dynamical evolution of NS and/or CBM created in heavy-ion collisions will be kept as a future excersize.  The paper is organized as follows. In section \\ref{sc.formalism} we discuss our model. In section \\ref{sc.results} we calculate different thermodynamic properties and present our result and finally we conclude in section \\ref{sc.conclusion}. ",
        "conclusions": "\\label{sc.conclusion}  In this paper we have studied the 2+1 flavor strongly interacting matter under the condition of $\\beta-$equilibrium. We have presented a comparative study of NJL versus PNJL model. The phase diagrams in these two models are broadly similar, but quantitatively somewhat different. The presence of the Polyakov loop delays the transition for larger values of temperature for a given quark chemical potential. As a result the CEP in the PNJL model is almost twice as hot as that in the NJL model. We have illustrated characteristics of the phase diagram with the behavior of some thermodynamic quantities like the constituent mass, compressibility, specific heat, speed of sound and the equation of state for $\\mu_e=0$ MeV and $\\mu_e=40$ MeV at $T$=50 MeV. We found striking differences between the NJL and PNJL model in terms of the softness of the equation of state in the hadronic and partonic phases.  The behavior of electric charge and baryon densities in the two models also differ in the hadronic phases to some extent. The differences become less with increasing electron density. We explained how the charge neutral trajectory is important in deciding the path along which the core of NS can change from hadronic to quark phase. For all values of $\\mu_e$ we find that the contours are all closed ones and give a restricted range of temperature and densities that are allowed. We speculated a possible scenario in which the quark-hadron transition in a NS would be a crossover. Again the baryon density contours seemed to suggest that if a system has baryon density three times the nuclear matter density it is quite surely in the partonic phase. We also found that the strangeness fraction increases steadily with increasing baryon density implying a possibility of having a strange NS.  The isentropic trajectories were obtained along which a system in hydrodynamic equilibrium is expected to evolve. The adiabats flow down from high temperature and low density towards low temperature and $\\mu_q=M_{vac}$, the constituent quark mass in vacuum. The adiabats then steeply rise along the transition line, thereafter goes towards higher densities with almost a constant slope. For small $s/n_B$ ratio the slope is so small that the isentropic trajectories almost become isothermal trajectories as well.  To summarize the scenario inside neutron stars we note that inside a newly born NS the temperature drops very quickly and gives rise to a system of low temperature nucleonic matter which may also be populated by hyperons and strange baryons due to high density near the core. The star is assumed to be $\\beta-$equilibrated and charge neutral. Now it is possible that due to some reason, {\\it e.g.} sudden spin down, this nucleonic matter will start getting converted to predominantly two flavor quark matter within strong interaction time scale. This transition would start at the center and a conversion front moving outward will convert much of the central region of the star. Along the path of the conversion front, each point inside the star may lie on an isentropic trajectory. Gradually this system of predominantly 2 flavor quark matter will get converted to strange quark matter through weak interactions and finally a $\\beta-$equilibrated charge neutral strange quark matter will be produced. The strangeness production occurs mainly through non-leptonic decay \\cite{ghosh-npa}, the system is expected to lie on a constant density line and move towards the point with highest strangeness possible at that density. Finally the semi-leptonic processes will take over and system will then evolve along a $\\beta-$equilibrated charge neutral contour.  The natural extension of the work is to obtain the detailed evolution of a family of neutron stars starting with different initial conditions and gravity effects incorporated. We hope to report the study in a future publication. It would also be important to consider colored exotic states like diquarks \\cite{bentz-npa} that may arise at high densities."
    },
    "1107/1107.4801_arXiv.txt": {
        "abstract": "We simulate the magnetic feature tracking (MFT) speed using advective-diffusive transport models in both one and two dimensions. By depositing magnetic bipolar regions at different latitudes at the Sun's surface and following their evolution for a prescribed meridional circulation and magnetic diffusivity profiles, we derive the MFT speed as a function of latitude. We find that in a one dimensional surface-transport model the simulated MFT speed at the surface is always the same as the meridional flow-speed used as input to the model, but is different in a two-dimensional transport model in the meridional ($r,\\theta$) plane. The difference depends on the value of the magnetic diffusivity and on the radial gradient of the latitudinal velocity. We have confirmed our results with two different codes in spherical and Cartesian coordinates. ",
        "introduction": "At the solar surface magnetic features are observed in the form of active regions. Tracking the motion of such a structure individually, one finds in general a poleward migration which is suggestive of a poleward meridional flow at and just beneath the Sun's surface. However, Doppler measurements of the poleward flow speed at the surface reveal a systematic difference from the speed inferred from magnetic feature-tracking (MFT): at low and mid latitudes, the latter is observed to be lower than the Doppler speed, but similar to it at high latitudes \\citep[see Fig.~10 of][]{ulrich2010}. Sunspots are usually discarded in such an analysis \\citep{komm+etal_93,hatha+righ_10} because their motion may be affected by their strong magnetic fields. To understand the physical origin of these differences, \\citet{dgu2010} performed a simple test using a 2D (axisymmetric) advective-diffusive flux-transport model. They showed in simulations  that, due to diffusive transport, the MFT speed can indeed be different from that of the meridional flow fed into the model. They attributed this difference to the latitudinal gradient of the  radial component of the magnetic field, directed towards the equator at the equatorward side of an active region and towards the pole at its poleward side. They concluded that magnetic features drift poleward with a net speed that is lower than the flow speed at low latitudes and higher at high latitudes.  In non-axisymmetric two-dimensional ($\\theta,\\phi$) surface-transport models \\citep[e.g.][]{baumann+etal_04,sheeley2010, wrs2009} one could likewise suppose that diffusion is the only agent that can prevent magnetic features from simply being advected with the meridional flow.  However, differences between Doppler and MFT speeds have never been discussed  for those models.  In this paper, we use both 1D and 2D advective-diffusive flux transport models to clarify to what extent the value  of the magnetic diffusivity and its radial gradient influence the difference between the circulation and MFT speeds. Moreover, we will study the role of the radial gradient of the latitudinal flow velocity. ",
        "conclusions": "Through 1D and 2D advective-diffusive models we have investigated the differences between the surface meridional flow speed obtained from Doppler measurements and that inferred from magnetic feature-tracking. In the one-dimensional simulations, the average velocity of the magnetic tracers always coincides with that of the flow, independently of the value of $\\etaT$. In 2D models, on the other hand, flow and feature-tracking velocities may diverge at higher diffusivities, for which the ``frozen-in'' condition does no longer hold. Further, the difference between these velocities depends on the radial gradient of the latitudinal velocity: the steeper $U_{\\theta}$, the larger the difference. Using a different code we have verified that these results apply also in Cartesian geometry. To understand this dependence we refer to the induction equation (for simplicity in Cartesian coordinates), taken at the surface  $x=L_x$ (or $r=R$) where $U_x=\\partial_y U_x=0$, so \\begin{alignat}{2} \\label{eq:bx} \\frac{\\partial B_x}{\\partial t}&=-&&\\frac{\\partial}{\\partial y} (U_yB_x) + \\etaT\\left(\\frac{\\partial^2 B_x}{\\partial x^2} + \\frac{\\partial^2 B_x}{\\partial y^2}  \\right),\\\\ \\label{eq:by} \\frac{\\partial B_y}{\\partial t}&=&&\\frac{\\partial}{\\partial x} (U_yB_x-U_xB_y)+\\etaT\\left(\\frac{\\partial^2 B_y}{\\partial x^2} + \\frac{\\partial^2 B_y}{\\partial y^2}  \\right). \\end{alignat} Note that $B_x$ is apparently decoupled from $B_y$, but at the price of being coupled to its second vertical derivative. In the case of small $\\etaT$ (i.e., diffusion time larger than advection time), the evolution of $B_x$ is governed by the first, advective, term in \\Eq{eq:bx}. In this case the magnetic field lines follow the fluid velocity locally (see curved magnetic field lines in the left hand panel of Fig. \\ref{fig:br1}). In the case of larger $\\etaT$ (diffusion time similar to or shorter than advection time), the diffusion term in \\Eq{eq:bx} plays a significant role in the field evolution. The dependence on $\\partial_x^2 B_x$ can be eliminated by the solenoidal condition $\\partial_x B_x+\\partial_y B_y =0$, by which the coupling between the two equations becomes obvious. Because $B_y$ depends explicitly on $\\partial_x U_y$, the surface speed of $B_x$ will clearly be modified by the fluid motion deeper down in the sub-surface layers. When ignoring $B_y$ from the beginning, however, as done in \\cite{devore+etal_84}, and many surface transport models afterwards, this influence will be lost. In spherical coordinates, the induction equation for $B_r$ exhibits an analogous dependence on $B_\\theta$, hence the same argument is valid.  Based upon our results for the difference between flow and feature-tracking speeds, one might think of inferring the thickness of the layer where the flow is poleward. This value, however, would depend on the surface diffusivity and on $\\partial_r U_{\\theta}$, both of which are poorly known. \\cite{hathaway_11} has inferred an extreme flow pattern with a very shallow poleward flow ($\\approx35\\Mm$ deep) which nevertheless can be brought into agreement with our findings, requiring a large radial  gradient of $U_\\theta$, that is, $\\bra{h_r(r)}\\ga 20$, cf. Fig.~\\ref{fig:veleta1w}. On the other hand, surface flux-transport models in two dimensions  ($\\theta,\\phi$) which disregard the radial derivatives in $B_r$, are probably overestimating the importance of advection in their results."
    },
    "1107/1107.4495_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec1} After the baryons are freed from the Compton drag of the photons\\footnote{In what follows the values of the redshifts $z$ will be estimated in terms of the standard $\\Lambda$CDM paradigm analyzed in the light of the WMAP 7yr data alone \\cite{wmap7a,wmap7b,wmap7c}.} (i.e. $z_{\\mathrm{drag}}  = 1020.3 \\pm 1.4$) large-scale magnetic fields contribute to the evolution of the inhomogeneities in the spatial curvature and in the matter distribution even if observational constraints on their effects are lacking.  After the drag epoch (i.e. $z < z_{\\mathrm{drag}}$) and before the Universe is reionized (i.e. $z \\geq z_{\\mathrm{re}}= 10.5 \\pm 1.2$) the Universe enters a very interesting epoch (often dubbed as dark age) when the first gravitationally bound system are formed. The purpose of the present paper is to explore the evolution of large-scale magnetic fields during the dark ages in the minimal framework of the magnetized $\\Lambda$CDM scenario  (m$\\Lambda$CDM in what follows) and in the simplest magnetohydrodynamical description which neglects the propagation of all the electromagnetic excitations of the plasma.  Probably the best evidence of reionization comes from the analysis of the large-scale temperature and polarization fluctuations in the Cosmic Microwave background (CMB). Roughly speaking WMAP finds that about $8$\\% of the CMB photons were scattered by free electrons in the intergalactic medium, but only $4$ \\% could have been scattered  by the intergalactic medium for redshifts $z< 6$. This means that the Universe must have been already reionized for $z \\geq {\\mathcal O}(10)$. The direct estimates of $z_{\\mathrm{re}}$ in terms of he WMAP 7yr data are consistent with  the Sloan Digital Sky Survey (SDSS) \\cite{sdss1,sdss2} reporting the observations of quasars whose absorption spectra show, for $z> 6$, a substantial increase in the fraction of ionized hydrogen. In \\cite{mg1,mg2} it has been argued, among other things, that the thickness of the last scattering surface and the optical depth at reionization are only mildly sensitive to the presence of an ambient magnetic field; the effect of reionization has then been parametrized, as in the conventional situation, with a second  Gaussian peak of the visibility function centered around the reionization time. Conversely the inhomogeneities of the geometry and of the sources are affected by the presence of the ambient magnetic field whose effects on the baryons and on the geometry can never be completely switched off.  Prior to photon decoupling the electrons, ions and photons are tightly coupled together; semi-analytic and numerical treatments of this problem are available in the presence of large-scale magnetic fields \\cite{mg1,mg2}. Useful applications of the analytic treatments contemplate, for instance, explicit formulas for the distortions and for the shifts in the positions of the acoustic peaks as well as of the polarization peaks. For $z < z_{\\mathrm{drag}}$ the evolution of the inhomonogeneities in the m$\\Lambda$CDM scenario simplifies in many respects and the evolution of the large-scale inhomogeneities can be solved exactly. The solutions for the evolution of the metric inhomogeneities in the m$\\Lambda$CDM paradigm can also be extended to complementary cases provided the dark energy does not cluster as it can be conveniently assumed for different applications. In the literature there exist some exact solutions of the evolution equations of metric perturbations in pressureless media like the ones of Hwang \\cite{hwang1,hwang2} or the ones discussed in Refs. \\cite{et1,et2,et3}. In the present paper after deriving the evolution equations of metric inhomogeneities for $z< z_{\\mathrm{drag}}$, the exact solutions of the perturbation equations will be reported in the presence of a fully inhomogeneous magnetic field in the magnetohydrodynamical approximation stipulating, as customarily done, that the magnetic field, the Ohmic current and the electric field are all solenoidal (comoving) vectors.  It is relevant to mention that various projects or radio arrays aim at a direct scrutiny of the dark ages like LOFAR (low-frequency array) \\cite{lofar}, MWA (Murchison Wide-Field Array) \\cite{mwa}, the enhanced VLA (Very Large Array) \\cite{VLA} and the SKA (square kilometer array) \\cite{SKA}. One of the common strategies of these programs is the direct observational tomography of the (redshifted) $21$ cm emission. The logic is to construct radio arrays capable of mapping the three-dimensional distribution of primeval hydrogen. Such a spatial distribution with its structures, its clumpiness, possibly its polarization bears the mark of the evolution through the dark ages and, most notably, of the first galaxies and stars producing ultraviolet emission and hence holes in the distribution of the protohydrogen. The role of cosmic magnetism in this epoch should be more carefully explored but the growth of matter inhomogeneities during the dark ages is certainly affected by the presence of large-scale magnetic fields, as general physical considerations suggest and as the solutions presented hereunder confirm. In section \\ref{sec2} we are going  to present the evolution of the inhomogeneities in the m$\\Lambda$CDM model and for $z< z_{\\mathrm{drag}}$. In section \\ref{sec3} the exact solutions of the system will be presented. In section \\ref{sec4}  the exact solutions are applied to the explicit calculation of the magnetized growth rate which is then compared to the conventional result of the concordance scenario with the purpose of deriving en explicit bound valid over the inhomogeneity scales relevant to structure formation.  \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": "\\label{sec4} The integral appearing in Eq. (\\ref{SR4}) can be evaluated explicitly and the obtained results expanded in the limit $\\alpha_{\\mathrm{i}}= a_{\\mathrm{i}}/a_{\\mathrm{eq}} > 1$ and $\\alpha= a/a_{\\mathrm{eq}} \\gg 1$ which corresponds to the physically interesting situation where the initial conditions of the solution (\\ref{SOLM3})  are set well after matter-radiation equality and anyway for redshifts smaller than $z_{\\mathrm{drag}}$. In the latter limit  the result can be expressed as \\begin{equation} \\delta_{\\mathrm{m}}^{(<)}(\\vec{x},\\alpha) = \\frac{2}{3}  \\biggl(\\frac{\\omega_{\\mathrm{b}0}}{\\omega_{\\mathrm{M}0}} \\biggr) \\frac{(z_{\\mathrm{eq}}+1)^3}{ H_{0}^2 \\Omega_{\\mathrm{M}0}} \\frac{{\\mathcal S}_{\\mathrm{B}}(\\vec{x},\\alpha)}{\\alpha^3} + {\\mathcal D}_{1}(\\vec{x}) \\biggl( 1 + \\frac{3}{2} \\alpha\\biggr)+ {\\mathcal O}\\biggl(\\frac{1}{\\alpha}\\biggr) + {\\mathcal O}\\biggl(\\frac{1}{\\alpha_{\\mathrm{i}}}\\biggr), \\label{SR8} \\end{equation} where, we remind, $\\alpha= a/a_{\\mathrm{eq}}$.  While in the regime $ a_{*} \\leq a \\ll a_{\\mathrm{eq}}$ the total matter density contrast is well estimated by Eq. (\\ref{SR8}),  for  $a > a_{*}$ the solution is instead given by Eq. (\\ref{SOLM3}). The initial conditions for the calculation of the growth rate can therefore be set by requiring that \\begin{equation} \\delta^{(<)}_{\\mathrm{m}}(\\vec{x}, a_{*}/a_{\\mathrm{eq}}) = \\delta_{\\mathrm{m}}(\\vec{x},a_{*}/a_{\\mathrm{de}}),\\qquad \\frac{\\partial}{\\partial a}  \\delta^{(<)}_{\\mathrm{m}}\\biggl|_{a = a_{*}} =  \\frac{\\partial}{\\partial a}  \\delta_{\\mathrm{m}}\\biggl|_{a = a_{*}}. \\label{C1} \\end{equation} The two conditions reported in Eq. (\\ref{C1}) fix the constants appearing in Eq. (\\ref{SOLM3}). The growth rate \\cite{gamma,growth} (see also \\cite{bb0}) can be computed from Eq. (\\ref{SOLM3}) and it is given by \\begin{equation} f(\\vec{x}, \\alpha) = \\frac{\\partial \\ln{\\delta_{\\mathrm{m}}}}{\\partial\\ln{\\alpha}}, \\label{GR1} \\end{equation} where, according to Eq. (\\ref{SOLM3}), $\\delta_{\\mathrm{m}}(\\vec{x},\\alpha)$ can be written as \\begin{equation} \\delta_{\\mathrm{m}}(\\vec{x},\\alpha) = \\delta(\\vec{x}) \\biggl[ {\\mathcal C}_{1}(w_{\\mathrm{de}}, z_{\\mathrm{de}}, z_{*}) {\\mathcal F}_{4}(\\alpha,w_{\\mathrm{de}}) + {\\mathcal C}_{2}(w_{\\mathrm{de}}, z_{\\mathrm{de}}, z_{*}) {\\mathcal F}_{5}(\\alpha,w_{\\mathrm{de}})\\biggr] + \\Sigma_{\\mathrm{B}}(\\vec{x}, \\alpha). \\label{GR2} \\end{equation} The two constants ${\\mathcal C}_{1}(w_{\\mathrm{de}}, z_{\\mathrm{de}}, z_{*})$ and ${\\mathcal C}_{2}(w_{\\mathrm{de}}, z_{\\mathrm{de}}, z_{*})$ are a lengthy combination of hypergeometric functions depending on $a_{*} = 1/(z_{*} +1)$ and $a_{\\mathrm{de}} = 1/(z_{\\mathrm{de}} +1)$. The quantity $\\delta(\\vec{x})$ denotes the spatial profile of the fluctuation of the total density contrast which can be estimated in terms of the matter power spectrum while $\\Sigma_{\\mathrm{B}}(\\vec{x},\\alpha)$ denotes the contribution of the magnetic sources. The power spectrum of $\\delta(\\vec{x})$ can be written as \\begin{eqnarray} \\langle \\delta(\\vec{x}) \\, \\delta(\\vec{x} + \\vec{r}) \\rangle &=& \\int d \\ln{k} \\, P_{\\delta}(k) \\, \\frac{\\sin{k\\,r}}{k\\, r}, \\nonumber\\\\ {\\mathcal P}_{\\delta}(k,y_{\\mathrm{eq}}) &=& \\frac{4}{25}{\\mathcal A}_{{\\mathcal R}} \\biggl(\\frac{k}{k_{\\mathrm{p}}}\\biggr)^{n_{\\mathrm{s}} -1} \\ln^2{(k/k_{\\mathrm{eq}})}, \\label{GR3} \\end{eqnarray} holding for wavenumbers larger than $k_{\\mathrm{eq}} = 0.00974_{-0.00040}^{0.00041}\\, \\mathrm{Mpc}^{-1}$. In Eq. (\\ref{GR3}) $k_{\\mathrm{p}}=0.002\\, \\mathrm{Mpc}^{-1}$ and ${\\mathcal A}_{{\\mathcal R}} = (2.43\\pm 0.11)\\times 10^{-9}$ in the case of  the WMAP 7yr data alone \\cite{wmap7a,wmap7b,wmap7c} and for the case of the $\\Lambda$CDM scenario.  Typical wavenumbers  $k \\gg k_{\\mathrm{eq}}$ crossed inside the Hubble volume before matter-radiation equality. Conversely the very large length-scales (relevant for the region of the Sachs-Wolfe plateau and for the integrated Sachs-Wolfe effect) fall into the complementary regime $k \\ll k_{\\mathrm{eq}}$. The wavenumbers  touched by the present considerations range from $k_{\\mathrm{min}} = 0.01 \\, h_{0}\\, \\mathrm{Mpc}^{-1}$ to, approximately, $k_{\\mathrm{max}} = 0.3 \\,h_{0}\\, \\mathrm{Mpc}^{-1}$. The range $k_{\\mathrm{min}} \\leq k \\leq k_{\\mathrm{max}}$ includes also the scale at which the spectrum becomes nonlinear, i.e. $k_{\\mathrm{nl}} \\geq 0.09 \\, h_{0} \\, \\mathrm{Mpc}^{-1}$.  The illustration of the analytical results will be given in terms of the following fiducial set of parameters determined on the basis of the WMAP 7yr data alone \\cite{wmap7a,wmap7b,wmap7c}: \\begin{equation} ( \\Omega_{\\mathrm{b}0}, \\, \\Omega_{\\mathrm{c}0}, \\Omega_{\\mathrm{de}0},\\, h_{0},\\,n_{\\mathrm{s}},\\, \\epsilon_{\\mathrm{re}}) \\equiv (0.0449,\\, 0.222,\\, 0.734,\\,0.710,\\, 0.963,\\,0.088), \\label{FL10} \\end{equation} where $\\epsilon_{\\mathrm{re}}$ denotes the optical depth at reionization and $n_{\\mathrm{s}}$ is the spectral index of curvature perturbations. The power spectra contributing to $\\Sigma_{\\mathrm{B}}(\\vec{x},\\alpha)$ can be solely expressed in  terms of $\\Omega_{\\mathrm{B}}(\\vec{x},\\alpha)$ and $\\sigma_{\\mathrm{B}}(\\vec{x},\\alpha)$ by recalling that \\begin{equation} \\frac{\\vec{\\nabla} \\cdot (\\vec{J} \\times \\vec{B})}{4 a^4 \\rho_{\\gamma}} =4 \\nabla^2 \\sigma_{\\mathrm{B}} - \\nabla^2 \\Omega_{\\mathrm{B}},\\qquad \\Omega_{\\mathrm{B}}(\\vec{x}, \\alpha) =\\frac{B^2(\\vec{x})}{8\\, \\pi \\, a^4 \\rho_{\\gamma}}, \\qquad \\partial_{i} \\partial_{j} \\Pi^{ij}_{\\mathrm{B}} = \\frac{4}{3} \\rho_{\\gamma} \\nabla^2 \\sigma_{\\mathrm{B}}, \\label{GR4} \\end{equation} where $\\Pi^{ij}_{\\mathrm{B}}$ denotes  the anisotropic stress of the magnetic fields. \\begin{figure}[!ht] \\centering \\includegraphics[height=5cm]{FIG1.eps} \\includegraphics[height=5cm]{FIG2.eps} \\caption[a]{The growth rate is illustrated for two different values of the barotropic index, i.e. $w_{\\mathrm{de}} = -1$ (full line) and $w_{\\mathrm{de}} =-0.6$ (dashed line). In the plot at the left $B_{\\mathrm{L}} =1 \\, \\mathrm{nG}$ and $k= 0.03\\, \\mathrm{Mpc}^{-1}$. In the plot at the right $B_{\\mathrm{L}} =10 \\, \\mathrm{nG}$ and $k= 0.01\\, \\mathrm{Mpc}^{-1}$. In both plots the magnetic spectral index has been chosen as $n_{\\mathrm{B}} = 1.5$.} \\label{F1} \\end{figure} In the range $k_{\\mathrm{min}} \\leq k \\leq k_{\\mathrm{max}}$, the terms containing the Ohmic current in the evolution equations (e.g. the term containing ${\\mathcal S}_{\\mathrm{B}}(\\vec{x},\\alpha)$ in ${\\mathcal T}^{(1)}_{\\mathrm{B}}(\\vec{x},\\alpha)$) dominates against the one containing just $\\Omega_{\\mathrm{B}}(\\vec{x},\\tau)$ because of the presence of two supplementary spatial gradients. The same comment can be made for $\\Sigma(\\vec{x},\\alpha)$. In Fig. \\ref{F1} the growth rate is reported for two illustrative values of the barotropic index $w_{\\mathrm{de}} =-1$ (full line) and $w_{\\mathrm{de}} = -0.6$ (dashed line); in both plots of Fig. \\ref{F1}, $k = 0.03\\, \\mathrm{Mpc}^{-1}$. In Fourier space the power spectra of $\\Omega_{\\mathrm{B}}$ and $\\sigma_{\\mathrm{B}}$ will be denoted by ${\\mathcal P}_{\\Omega}(k)$ and ${\\mathcal P}_{\\sigma}(k)$  and their explicit expression can be written as \\begin{equation} {\\mathcal P}_{\\Omega}(k) = {\\mathcal E}_{\\mathrm{B}} \\biggl(\\frac{k}{k_{\\mathrm{L}}}\\biggr)^{2 (n_{\\mathrm{B}} -1) +\\alpha_{\\Omega}},\\qquad {\\mathcal P}_{\\sigma}(k) = r_{\\mathrm{B}} \\,{\\mathcal E}_{\\mathrm{B}} \\biggl(\\frac{k}{k_{\\mathrm{L}}}\\biggr)^{2 (n_{\\mathrm{B}} -1) +\\alpha_{\\sigma}}, \\label{PS1} \\end{equation} $\\alpha_{\\Omega}$ and $\\alpha_{\\sigma}$ are the corresponding running of the spectral indices which are set to zero in the minimal magnetized $\\Lambda$CDM scenario. The relative amplitude of the two power spectra at the magnetic pivot scale $k_{\\mathrm{L}}$ (conventionally chosen to be $1\\,\\, \\mathrm{Mpc}^{-1}$) is controlled by $r_{\\mathrm{B}}$. Since the magnetic fields are, themselves, stochastically distributed,  the ensemble average of their Fourier modes obeys \\begin{equation} \\langle B_{i}(\\vec{k}) \\, B_{j}(\\vec{p}) \\rangle = \\frac{2\\pi^2 }{k^3} P_{ij}(k) {\\mathcal P}_{\\mathrm{B}}(k) \\delta^{(3)}(\\vec{k} + \\vec{p}), \\qquad {\\mathcal P}_{{\\mathrm{B}}}(k) = {\\mathcal A}_{\\mathrm{B}} \\biggl(\\frac{k}{k_{\\mathrm{L}}}\\biggr)^{n_{\\mathrm{B}}-1}, \\label{PS2} \\end{equation} where $P_{ij}(k) = (k^2 \\delta_{ij} - k_{i} k_{j})/k^2$ and  ${\\mathcal A}_{\\mathrm{B}}$ the spectral amplitude of the magnetic field. The spectral amplitude of the magnetic energy density ${\\mathcal E}_{\\mathrm{B}}$ depends upon the ultraviolet cut-off associated with diffusion damping (i.e. $k_{\\mathrm{D}}$), upon the infrared cut-off associate with the (comoving) angular diameter distance at last scattering (i.e. $D_{\\mathrm{A}}(z_{1})$) and also on the magnetic spectral index $n_{\\mathrm{B}}$, i.e. \\cite{mg1,mg2} \\begin{eqnarray} && r_{\\mathrm{B}} = \\frac{n_{\\mathrm{B}} + 29}{20 (7 - n_{\\mathrm{B}})}\\biggl[ 1 + \\frac{(5 - 2 n_{\\mathrm{B}})(83 n_{\\mathrm{B}}- 473)}{2 ( 7 - n_{\\mathrm{B}}) (n_{\\mathrm{B}} + 29)} \\biggl(\\frac{k}{k_{\\mathrm{A}}}\\biggr)^{1 - n_{\\mathrm{B}}}\\biggr], \\qquad n_{\\mathrm{B}} < 1, \\label{PS3}\\\\ && r_{\\mathrm{B}}  = \\frac{n_{\\mathrm{B}} + 29}{20 (7 - n_{\\mathrm{B}})}\\biggl[1 + \\frac{(n_{\\mathrm{B}} -1)(87 n_{\\mathrm{B}} -501)}{(n_{\\mathrm{B}} + 29) ( 7 - n_{\\mathrm{B}})}\\biggl(\\frac{k}{k_{\\mathrm{D}}}\\biggr)^{5 - 2 n_{\\mathrm{B}}}\\biggr] ,\\qquad n_{\\mathrm{B}} >1, \\label{PS4} \\end{eqnarray} where, recalling that $D_{A}(z_{1}) = 2 d_{\\mathrm{A}}(z_{1})/(H_{0} \\sqrt{\\Omega_{\\mathrm{M}0}})$, \\begin{equation} k_{\\mathrm{A}}(z_{1}) = 1/D_{\\mathrm{A}}(z_{1}),\\qquad \\frac{k_{\\mathrm{D}}(z_{1})}{k_{\\mathrm{A}}(z_{1})}= \\frac{2240 \\, d_{\\mathrm{A}}(z_{1})}{\\sqrt{\\sqrt{r_{\\mathrm{R}1} +1} - \\sqrt{r_{\\mathrm{R}1}}}} \\biggl(\\frac{z_{1}}{10^{3}} \\biggr)^{5/4} \\, \\omega_{\\mathrm{b}}^{0.24} \\omega_{\\mathrm{M}}^{-0.11}. \\label{PS4a} \\end{equation} In Eq. (\\ref{PS4a}) $r_{\\mathrm{R}1}$ is the ratio of radiation to matter at $z_{1}$ which is the redshift of last-scattering and which can be determined analytically for typical values of the parameters close to the best-fit determined on the basis of the WMAP 7yr data and in the light of the $\\Lambda$CDM paradigm: \\begin{eqnarray} && r_{\\mathrm{R}1} = \\frac{\\rho_{\\mathrm{R}}(z_{1})}{\\rho_{\\mathrm{M}}(z_{1})} = \\frac{ 4.15\\times 10^{-2}}{\\omega_{\\mathrm{M}}}\\, \\biggl(\\frac{z_{1}}{10^{3}}\\biggr), \\qquad z_{1} = 1048[ 1 + f_{1} \\omega_{\\mathrm{b}}^{- f_{2}}] [ 1 + g_{1} \\omega_{\\mathrm{M}}^{g_2}], \\nonumber\\\\ && g_{1} = \\frac{0.0783 (\\omega_{\\mathrm{b}})^{-0.238}}{[1 + 39.5 (\\omega_{\\mathrm{b}})^{0.763}]},\\qquad g_{2} = \\frac{0.560}{1 + 21.1 (\\omega_{\\mathrm{b}})^{1.81}}, \\label{PS4b} \\end{eqnarray} where $f_{1}$ and $f_{2}$ are given, respectively, by $f_{1} = 1.24\\times 10^{-3}$, $f_{2} = 0.738$.  When $n_{\\mathrm{B}}=1$ the spectra for the energy density and for the Lorentz force are scale-invariant once the logarithmic divergence of the two-point function is appropriately subtracted.  The amplitude ${\\mathcal A}_{\\mathrm{B}}$ can be traded for the magnetic field regularized over a spatial domain $k_{\\mathrm{L}}^{-1}$, i.e. in the case $n_{\\mathrm{B}} > 1$ \\begin{equation} {\\mathcal E}_{\\mathrm{B}} = \\frac{4 ( 7 - n_{\\mathrm{B}})}{3 (n_{\\mathrm{B}} -1) (5 - 2 n_{\\mathrm{B}})} \\frac{{\\mathcal A}_{\\mathrm{B}}^2}{(8\\pi \\overline{\\rho}_{\\gamma})^2}, \\qquad {\\mathcal A}_{\\mathrm{B}} = \\frac{(2\\pi)^{n_{\\mathrm{B}} -1}}{\\Gamma[(n _{\\mathrm{B}}-1)/2]} B_{\\mathrm{L}}^2, \\label{PS5} \\end{equation} and analogously in the $n_{\\mathrm{B}} <1$ case but with ${\\mathcal A}_{\\mathrm{B}} =[ (1 -n_{\\mathrm{B}})/2] (k_{\\mathrm{A}}/k_{\\mathrm{L}})^{(1 - n_{\\mathrm{B}})}B_{\\mathrm{L}}^2$.  In the case of the minimal m$\\Lambda$CDM the relation of $(n_{\\mathrm{B}}, {\\mathcal E}_{\\mathrm{B}})$ to $(n_{\\mathrm{B}}, B_{\\mathrm{L}})$ follows from Eqs. (\\ref{PS2})--(\\ref{PS5}). The magnetic energy density can be naturally referred to the energy density of the photon background so that $\\Omega_{\\mathrm{B\\,L}} = B_{\\mathrm{L}}^2/(8\\pi \\overline{\\rho}_{\\gamma})$ can be measured in units of the amplitude of curvature perturbations: \\begin{equation} \\frac{\\overline{\\Omega}_{\\mathrm{B\\,L}}}{{\\mathcal A}_{\\mathcal R}}=  39.568 \\biggl(\\frac{B_{\\mathrm{L}}}{\\mathrm{nG}}\\biggr)^{2} \\, \\biggl(\\frac{T_{\\gamma0}}{2.725\\,\\mathrm{K}}\\biggr)^{-4} \\, \\biggl(\\frac{{\\mathcal A}_{{\\mathcal R}}}{2.41\\times 10^{-9}}\\biggr)^{-1}. \\label{PS6} \\end{equation} The parameter space  of the magnetized $w$CDM models and of the magnetized  $\\Lambda$CDM scenario have been investigated, respectively, in \\cite{mg1} and \\cite{mg2}. In a frequentist approach, the boundaries of the confidence regions obtained in \\cite{mg1,mg2} represent exclusion plots at  $68.3\\,$\\% and $95.4\\,$\\% confidence level. When moving from the magnetized $\\Lambda$CDM scenario to the magnetized $w$CDM model we have that the the parameters maximizing the likelihood get shifted to slightly larger values\\footnote{The difference between Eqs. (\\ref{NTT}) and (\\ref{NTE}) is that the parameters of Eq. (\\ref{NTT}) are obtained from the analysis of the temperature autocorrelations while the parameters of Eq. (\\ref{NTE})  are obtained by adding the data points of the cross-correlations between temperature and E-mode polarization \\cite{mg1}.} \\begin{eqnarray} &&(n_{\\mathrm{B}},B_{\\mathrm{L}})_{\\Lambda\\mathrm{CDM}} = ( 1.598,\\,3.156 \\mathrm{nG})\\to (n_{\\mathrm{B}} , B_{\\mathrm{L}})_{\\mathrm{{\\it w}CDM}} =(1.883, \\,4.982\\, \\mathrm{nG}), \\label{NTT}\\\\ &&(n_{\\mathrm{B}}, B_{\\mathrm{L}})_{\\Lambda\\mathrm{CDM}} = ( 1.616,\\,3.218 \\mathrm{nG})\\to (n_{\\mathrm{B}},B_{\\mathrm{L}})_{\\mathrm{{\\it w}CDM}} = (1.913, \\,5.163\\, \\mathrm{nG}). \\label{NTE} \\end{eqnarray} Even if the addition of a fluctuating dark energy background pins down systematically larger values of the magnetic field  parameters, the results of \\cite{mg1,mg2} will be used here just for a consistent illustration of the results. \\begin{figure}[!ht] \\centering \\includegraphics[height=5cm]{FIG3.eps} \\includegraphics[height=5cm]{FIG4.eps} \\caption[a]{In the plot at the left the growth rate is illustrated for $\\alpha =1$ (i.e. $a= a_{\\mathrm{de}}$) for $B_{\\mathrm{L}} = 1\\, \\mathrm{nG}$ (full line) and $B_{\\mathrm{L}} = 10 \\, \\mathrm{nG}$ (dashed line). In the right plot $\\alpha =0.1$ (i.e. $a = 0.1 \\, a_{\\mathrm{de}}$) and the values of $B_{\\mathrm{L}}$ are, respectively, $B_{\\mathrm{L}} = 10\\, \\mathrm{nG}$ (full line) and $B_{\\mathrm{L}} = 30\\, \\mathrm{nG}$ (dashed line). In both plots $n_{\\mathrm{B}} =1.5.$ and $k =0.03\\, \\mathrm{Mpc}^{-1}$.} \\label{F2} \\end{figure} In Fig. \\ref{F2} the growth rate is illustrated as a function of $w_{\\mathrm{de}}$ and for two different choices of $B_{\\mathrm{L}}$, i.e. $1\\, \\mathrm{nG}$ (full line) and $10\\, \\mathrm{nG}$ (dashed line).  The difference between the plot at the left and at the right is the value of the redshift: while in the plot at the left $a= a_{\\mathrm{de}}$ (i.e. $z = z_{\\mathrm{de}}$)  in the plot at the right $a= 0.1\\,a_{\\mathrm{de}}$ and $z > z_{\\mathrm{de}}$. \\begin{figure}[!ht] \\centering \\includegraphics[height=5cm]{FIG5.eps} \\includegraphics[height=5cm]{FIG6.eps} \\caption[a]{In both plots the growth rate is illustrated for $a= a_{\\mathrm{eq}}$ as a function of the regularized magnetic field intensity $B_{\\mathrm{L}}$. In the plot at the left $w_{\\mathrm{de}} = -1$ (and $k =0.03\\, \\mathrm{Mpc}^{-1}$) while in the plot at the right $w_{\\mathrm{de}} = -0.6$ (and $k=0.3\\,\\mathrm{Mpc}^{-1}$).} \\label{F3} \\end{figure} In Fig. \\ref{F3} the growth rate is illustrated as a function of $B_{\\mathrm{L}}$ for different choices of the wavenumber and of the barotropic index. The full and dashed lines correspond, in each plot of Fig. \\ref{F3}, to two different values of the magnetic spectral index, i.e. $n_{\\mathrm{B}} = 1.5$ (full line) and $n_{\\mathrm{B}} =1.2$ dashed line.  In the absence of large-scale magnetic fields, the growth rate can be parametrized as $\\Omega_{\\mathrm{M}}^{\\gamma}(\\alpha)$ where $\\gamma$ is the growth index which can be explicitly estimated as \\cite{gong} (see also \\cite{gamma}) \\begin{equation} \\gamma(w_{\\mathrm{de}}) =\\frac{ 3 w_{\\mathrm{de}} - 3}{6 w_{\\mathrm{de}} - 5} + \\frac{3}{125} \\frac{(1- w_{\\mathrm{de}})(1 - 3 w_{\\mathrm{de}}/2)}{(1- 6 w_{\\mathrm{de}}/5)^2}\\epsilon +  {\\mathcal O}(\\epsilon^2), \\label{BB1} \\end{equation} where $\\epsilon = 1 - \\Omega_{\\mathrm{M}}$. The requirement  that the magnetized growth rate of Eq. (\\ref{GR1}) does not exceed the standard fit for the growth rate implies an upper bound on $B_{\\mathrm{L}}$ which is illustrated in Fig. \\ref{F4}. \\begin{figure}[!ht] \\centering \\includegraphics[height=5cm]{FIG7.eps} \\includegraphics[height=5cm]{FIG8.eps} \\caption[a]{The bound on $B_{\\mathrm{L}}$ illustrated for $n_{\\mathrm{B}} = 1.6$ (long dashed line), $n_{\\mathrm{B}} = 1.5$ (full line) and $n_{\\mathrm{B}} = 1.4$ (short dashed line). In the left plot $k=0.03\\, \\mathrm{Mpc}^{-1}$. In the right plot $k=0.3\\, \\mathrm{Mpc}^{-1}$. In both plots $a = a_{\\mathrm{eq}}$} \\label{F4} \\end{figure} The bound on $B_{\\mathrm{L}}$ gets more stringent as $k$ increases and for large redshift. In Fig. \\ref{F4} $a= a_{\\mathrm{de}}$. The allowed values of $B_{\\mathrm{L}}$ for different values of $w_{\\mathrm{de}}$ stay below each of the curves reported in the plots of Fig. \\ref{F4}. Recalling the definition of the comoving magnetic Jeans length \\cite{bb0} \\begin{equation} \\lambda_{\\mathrm{B\\,J}} = c_{\\mathrm{a}} \\sqrt{\\frac{\\pi}{G \\, a^3 \\rho_{\\mathrm{b}}}}=1.90\\times 10^{-2} \\, \\biggl(\\frac{\\omega_{\\mathrm{b}0}}{0.02258}\\biggr)^{-1} \\, \\biggl(\\frac{B_{\\mathrm{L}}}{\\mathrm{nG}}\\biggr)\\,\\,\\mathrm{Mpc}, \\qquad c_{\\mathrm{a}}^2 = \\frac{B_{\\mathrm{L}}^2}{8 \\pi a^3 \\,\\rho_{\\mathrm{b}}}, \\label{MJ1} \\end{equation} the results of Fig. \\ref{F4} imply $\\lambda_{\\mathrm{B\\, J}}  \\leq {\\mathcal O}(\\mathrm{Mpc})$. The considerations reported in the present analysis suggest that the study of large-scale magnetism during the dark ages provide further constraints on the evolution of large-scale magnetic fields. The quantitative features of the obtained constraints have been shown to be both complementary and competitive with the bounds stemming from a direct analysis of potential distortions induced on the temperature and polarization anisotropies for a range of length scales shorter than the one relevant for the physics of the CMB \\cite{rev}.  If the origin of the large-scale magnetic fields is primordial (as opposed to astrophysical) it is plausible to expect the presence of magnetic fields in the primeval plasma also  before the decoupling of radiation from matter.  If the evolution of the large-scale magnetic fields follows the tenets of magnetohydrodynamics then the magnetic flux will be conserved across the last scattering and potential effects of the magnetic fields during the dark ages could help to decide on their origin and implications.   In this perspective the SKA \\cite{SKA} program could  provide decisive informations both from the foreseen full sky surveys of Faraday rotation and  also from more detailed pictures of structure formation and reionization over a wide range of redshifts during the dark age.  \\newpage"
    },
    "1107/1107.0669_arXiv.txt": {
        "abstract": "We discuss the response of  neutron stars to the tidal interaction in a compact binary system, as encoded in the Love number associated with the induced deformation. This problem is of interest for gravitational-wave astronomy as there may be a detectable imprint on the signal from the late stages of  binary coalescence. Previous work has focussed on simple barotropic neutron star models, providing an understanding of the role of the stellar compactness and overall density profile. We add realism to the discussion by developing the framework required to model stars with varying composition and an elastic crust. These effects are not expected to be significant for the next generation of detectors but it is nevertheless useful to be able to quantify them. Our results show that (perhaps surprisingly) internal stratification has no impact whatsoever on the Love number. We also show that crust elasticity provides a (predictably) small correction to existing models. ",
        "introduction": "\\label{sec:intro}  When a massive body is exposed to a relatively weak external gravitational field we expect it to respond by changing shape. This is most easily understood by considering the gravitational effects of the Moon on the Earth. The oceans move to reach equilibrium as the moon orbits, leading to the observed tides. The effect will also deform the Earth's elastic crust, again to reach an equilibrium with the passing body. The deformation of the massive body modifies the tidal potential, an effect that can be expressed in terms of dimensionless quantities known as the Love numbers. We will consider this problem for a neutron star in a close binary system, focussing on the tidal Love number; the measure of the tidal response of an object to an external gravitational field \\cite{Stacey}.  Specifically, we are interested in $k_2$, the measure of a tidal deformation by a quadrupole perturbation due to an external gravitational field. This quantity is of direct relevance for future gravitational-wave astronomy, as it provides a potentially detectable ``finite-size correction'' to the late stages of a binary inspiral signal \\cite{Flanagan:2008}.  The effect depends on the density distribution of the star, and hence may be used as a diagnostic to use observational data to constrain neutron star theory. The implications for observations, and the formalism required to work out the quadrupole deformation in relativity, were first discussed in  \\cite{Flanagan:2008}. Since then the formalism has been generalized to determine, in particular, the shape Love number, see for example \\cite{Poisson:2009,Damour:2009}. Applications have mainly considered barotropic fluid models, with realistic equations of state studied in \\cite{Damour:2009} and the difference between neutron stars and self-bound quark stars being quantified in~\\cite{Postnikov:2010}.  The main aim of the present work is to develop the formalism required to study the tidal deformation of real neutron stars, accounting for key aspects like the elastic crust and variations in the interior composition. These are important developments, even though we do not expect these features to be easily observed in a binary gravitational-wave signal. The Love number leaves an imprint that is borderline detectable by advanced LIGO \\cite{Flanagan:2008}. The effect will be more important for the third generation Einstein Telescope \\cite{READ}, but small corrections to it may not seem that relevant. Nevertheless, as a matter of principle it is important to  move beyond back-of-the-envelope estimates and quantify the actual effect. Moreover, developments in this direction are required to address a number of questions of more immediate astrophysical relevance. By developing the required relativistic perturbation formalism for the static deformations induced by the tidal interaction, we prepare the ground for studies of general crust asymmetries, e.g., neutron star ``mountains'' which lead to gravitational-wave emission at twice the star's spin frequency \\cite{Jaranowski:1998,LIGO}. The work presented here represents key steps towards modelling such mountains in general relativity, which would allow us to consider realistic equations of state for the first time in this context. A closely related problem to that considered here concerns the breaking of the crust during binary inspiral. It is easy to estimate (based on the energetics involved~\\cite{Postnikov:2010}) that the crust will break before the final merger, but when exactly does this happen and where does the crust yield first? Another related problem concerns the quasiperiodic oscillations observed in the tails of magnetar flares \\cite{Israel,Watts:2007}. Again, a key issue concerns the build-up of stresses in the crust of these strongly magnetized stars, the eventual fracture and the induced dynamics. The formalism developed here provides a useful first step for a discussion of this problem as well. ",
        "conclusions": "\\label{sec:conclusion}   We have extended the discussion of tidally deformed relativistic stars by including the effects of internal composition stratification and the presence of an elastic crust,  relevant for mature neutron stars. The most important development concerns the formalism required to account for more realistic neutron star physics. Building on a recent extension \\cite{RELLAG} of relativistic Lagrangian perturbation theory to the multi-fluid setting (allowing for one of the components to be elastic), we have formulated the tidal deformation problem for realistic neutron star models. The final model is obviously still ``incomplete\", as we are ignoring the magnetic field and we are not accounting for superfluidity  (which should be relevant both in the crust and the fluid core), but it nevertheless allows us to consider astrophysically relevant questions.  This paper should be seen as a first, proof-of-principle, study of the relevant issues. We have shown that (perhaps unexpectedly) the tidal deformations are not affected at all by composition variations in the star's core. Having implemented the formalism required to account for the crust elasticity, and solved the problem numerically for a simple model problem (based on a polytropic equation of state and a simplistic model for the crust shear modulus), we have also demonstrated that the presence of the crust has a (predictably) small effect on the tidal Love number. We have considered how this effect varies with the crust parameters, and confirmed that the results agree with intuition.  From a formal point-of-view, these are important developments, but it should be stressed that we do not expect the new features to leave an observable imprint on a binary gravitational-wave signal. This is fairly obvious since the Love number leaves an imprint that is barely detectable by future generations of gravitational-wave detectors in the first place \\cite{Flanagan:2008} and \\cite{READ}, and we are quantifying ``small'' corrections to it. Having said that, it is clear that we need to move beyond  back-of-the-envelope estimates in this problem area and develop the computational technology required to model real neutron stars and actual astrophysical scenarios.  The present work represents an important step towards this goal. By developing the required relativistic perturbation formalism for the tidal interaction problem, we lay the foundation for work on general crust asymmetries, e.g., neutron star ``mountains''  relevant for  gravitational-wave astronomy.  This problem has not yet been considered in general relativity, which is required if we want to use realistic equations of state. A closely related problem concerns the fracture of the crust during binary inspiral, the effect that this may have on a detected gravitational-wave signal and possible associated electromagnetic signatures. Another important problem concerns the quasiperiodic oscillations observed in the tails of magnetar flares \\cite{Israel, Watts:2007,Samuelsson:2009}. In this context, a key issue concerns the build-up of stresses in the crust of these strongly magnetized stars, the eventual fracture and the induced dynamics of the coupled magneto-elastic system.  These are exciting problem of immediate astrophysical interest, and we hope to make progress on them in the near future."
    },
    "1107/1107.5750_arXiv.txt": {
        "abstract": "This paper reports the discovery and characterization of the transiting hot giant exoplanet \\koib. The planet has an orbital period of $\\koicurLCPshort$ days, and radial velocity measurements from the {\\it  Hobby-Eberly Telescope} show a Doppler signal of $\\koicurRVK$ $\\ms$. From a transit-based estimate of the host star's mean density, combined with an estimate of the stellar effective temperature \\teff=$\\koicurSMEteff$ from high-resolution spectra, we infer a stellar host mass of \\koicurYYmlong ~\\msun\\ and a stellar radius of \\koicurYYrlong ~\\rsun. We estimate the planet mass and radius to be  $\\mpl = \\koicurPPmlong\\,\\mjup$ and $\\rpl = \\koicurPPrlong\\,\\rjup$. The host star is active, with dark spots that are frequently occulted by the planet. The continuous monitoring of the star reveals a stellar rotation period of 11.89 days, 8 times the the planet's orbital period; this period ratio produces \\emph{stroboscopic} effects on the occulted starspots. The temporal pattern of these spot-crossing events shows that the planet's orbit is prograde and the star's obliquity is smaller than 15$^\\circ$. We detected planetary occultations of \\koib\\ with both the {\\it Kepler} and \\spitzer\\ {\\it Space Telescopes}. We use these observations to constrain the eccentricity, $e$, and find that it is consistent with a circular orbit ($e<0.011$). The brightness temperatures of the planet the infrared bandpasses are T$_{\\rm{3.6~\\micron}}$=$1880\\pm100$~K and T$_{\\rm{4.5~\\micron}}$=$1770\\pm150$~K. We measure the optical geometric albedo $A_g$ in the \\kepler\\ bandpass and find $A_g = 0.10 \\pm 0.02$. The observations are best described by atmospheric models for which most of the incident energy is re-radiated away from the day side. ",
        "introduction": "NASA's \\kepler\\ mission is a space-based photometric telescope dedicated to finding Earth-size planets in the habitable zones of their host stars and to determining the frequency and characteristics of planetary systems around Sun-like stars \\citep{borucki10}. While monitoring nearly continuously the brightness of about 150,000 dwarf stars since May 2009 to achieve the above goals, it accumulates an extraordinary volume of data allowing atmospheric studies of the transiting hot-Jupiters present in the field of view \\citep{borucki09}. The first four months of observations were released and 1235 transiting planet candidates were reported \\citep{borucki11}. Because of their short periods and relatively deep transit depths, several hot-Jupiter candidates were discovered early on during the first weeks of the mission. Amongst those, the {\\it Kepler Object of Interest} KOI-203 was identified as a promising target, was selected for follow-up studies, and placed on the short-cadence target list for subsequent quarters. This is the object of interest of the current study.  The highly irradiated transiting hot-Jupiters currently provide the best opportunities for studying exoplanetary atmospheres in emission, during planetary occultations, when the exoplanets pass behind their parent stars \\citep{seager00,sudarsky00,fortney05,barman05}. Light emitted from exoplanets was first detected from space at infrared wavelengths \\citep{charbonneau05,deming05} and more recently in the optical \\citep{alonso09a,snellen09,borucki09,alonso10} using the \\corot\\ and \\kepler\\ {\\it Space Telescopes}. Notably, \\cite{rowe06,rowe08} used the {\\it Microvariablity and Oscillations of Stars} ({\\it MOST}) telescope to place a very stringent upper limit on the depth of the occultation of \\hda.   The hot-Jupiters detected by \\corot\\ and \\kepler\\ are particularly good targets for studying planetary atmospheres, because the precise and nearly uninterrupted photometric surveillance provided by the space satellites can allow the planetary occultations to be detected at optical wavelengths \\citep{snellen09,borucki09,desert11b}. Obtaining multiple wavelength observations of the relative depths of planetary occultations is necessary to constrain the broad band emergent spectra. Such observations are fundamental to understanding the energy budget of these objects \\citep{sudarsky03,burrows05,burrows08,spiegel10} and for comparative exoplanetology.    We confirm here the planetary status of KOI-203, designating it as \\koib, and study its atmosphere. We first describe the \\kepler\\ observations and transit modeling as well as the follow-up observations used to confirm this planet, including an orbital solution using radial velocities obtained with the {\\it High Resolution Spectrometer} (HRS) at the {\\it Hobby-Eberly Telescope} (HET). We then study the impact of the stellar variability on the transit lightcurves and we use occulted dark starspots to constrain the stellar obliquity. We finally combine occultation measurements obtained in the optical with \\kepler\\ and in the infrared with \\spitzer\\ to learn about the atmospheric properties of \\koib.       We first describe the observations, time series and analysis of the \\kepler\\ photometry in Sect.~2, then we describe the follow-up observations that confirm the planet detection in Sect.~3 and then discuss the stellar properties in Sect.~4. In Sect.~5, we describe the occultation measurements obtained from the visible and infrared. We present a global Monte-Carlo analysis of the complete sample of observations in Sect.~6 and finally discuss our findings in Sect.~7. ",
        "conclusions": "\\label{results}    \\subsection{Characteristics of the system \\koi\\ } \\label{sec:system}   Adopting the stellar mass and radius for \\koi\\ that we found in Section~\\ref{sec:stellar}, we obtain a mass of $\\mpl = \\koicurPPmlong\\,\\mjup$ and a radius of $\\rpl = \\koicurPPrlong\\,\\rjup$ for the planet, which leads to a density of $\\rhopl = \\koicurPPrho\\,\\gcmc$, and a surface gravity $\\loggpl = \\koicurPPlogg $ (cgs). We note that the surface gravity can be derived from the photometry and the velocimetry only \\citep{southworth07}. The position of \\koib\\ on the mass/radius diagram is fairly common and it appears to be slightly inflated compared to models that include the effects of stellar irradiation (e.g. \\citealt{latham10}).  The current upper limit on the orbital eccentricity $e$ from radial velocity measurements is consistent with zero for \\koib. We also measure the mid-occultation timing offset from both \\kepler\\ and \\spitzer\\ observations. The determination of the timing of the secondary eclipse constrains the planet's orbital eccentricity. Our estimate for the best-fit timing offset translates to a constraint on $e$ and the argument of pericenter $\\omega$. The timing is used to constrain the $e\\cos(\\omega)$ and we find that $e < 0.011$ at the $1-\\sigma$ level. This upper limit implies that the orbits of these objects are nearly circular unless the line of sight is aligned with the planet's major orbital axis, i.e. the argument of periapse $\\omega$ is close to 90\\degr~or 270\\degr.    We fit a linear function of transit epoch and find $T_{c}(0)=\\koicurLCT$ (BJD$_{utc}$) with a period $P=\\koicurLCP$ days. We also fit each mid-transit time individually and compare each one to the expected linear ephemeris to obtain a Observed-minus-Computed (O-C) time series. The O-C times shows a scatter that is most probably caused by stellar spot-induced shifts, since the star is active (see Section~\\ref{starvar}). Formally, the individual fits are consistent with the linear model, therefore we do not consider this to be a significant detection of timing anomalies. Since there is no clear evidence for transit timing variations (TTVs; \\citealt{agol05,holman05}), we use the timing data to place upper limits on the mass of a hypothetical second planet that would perturb the orbit of the transiting planet using the procedure described in \\cite{carter11}. Figure~\\ref{fig:ttvs} shows the constraints on the perturber mass as a function of period ratio, as determined from this analysis. The mass constraints on the perturber are more restrictive near the mean-motion resonances and most restrictive at the low-order resonances, particularly for the interior and exterior 2:1 resonances. For example, a perturber at the interior 2:1 resonance having a mass near that of Mars would have induced TTVs detectable in the present data.   \\subsection{Stellar variability of \\koi\\ } \\label{starvar} The \\kepler\\ photometry exhibits a quasi periodic flux modulation of about $3\\%$ with a period of  11.9~days (Figure~\\ref{fig:rawtlc}). This period corresponds to 8 times the planet's orbital period (Figure~\\ref{fig:stellarcycle}). The stellar variability can be quantified from a Lomb-Scargle periodogram that reveals two peaks at 5.95 and 11.9 days (see Figure~\\ref{fig:ls}). The position of the peak on the periodogram and its width provide a measurement of the stellar rotation period and its error, respectively. We measured a stelllar rotation period of $11.9\\pm1.1$~days. The presence of active stellar spots localized on opposite hemispheres can best explain the two peaks seen in the periodogram. We find that the ratio between the planet's orbital period and the stellar rotation period is  $8.0\\pm0.7$~days. We note that this intriguing integer ratio of 8 could potentially reveal the signature of stellar-planet interactions.       \\subsubsection{Using occulted starspots to infer the orbital obliquity} \\label{starspots}  The lightcurve of \\koi\\ shows substantial deformation of the planet's transit profiles which we interpret as occultations of dark starspots (Figure~\\ref{fig:keplerlc}). We co-added the transit lightcurves that occurred at epochs modulo the stellar rotation period (modulo 8 planetary orbital periods). We present the resulting 8 transit lightcurves in Figure~\\ref{fig:spotseq}; each contains about 22 individual transit lightcurves. We identified that the planet crosses spots of identical shapes every 8 transits producing a \\emph{stroboscopic} effect. Therefore, we conclude that starspots come back to the same position on the transit chord after 8 orbits or 11.89 days (one stellar rotation period). Because of this stroboscopic effect, the scatter measured during the spot-crossing phase of each of these lightcurve residuals is similar to the scatter of the residuals measured outside the transit. This further supports the assertion that the same spots are crossed by the planet every 8 transits, at identical orbital phases, because the residuals measured from a transit light curve obtained from a planet occulting a random distribution of spots are expected to be larger during the transit event compared to the out-of-transit monitoring. This can be explained only if spots are coming back to similar longitudes and latitudes after one stellar rotation period. Interestingly, occulted stellar spots can place constraints on the spin-orbit alignment  (e.g \\citealt{deming11,nutzman11,sanchis11a,sanchis11b}).   We identified five spots (A, B, C, D and E) by their anomaly in the residuals and we marked their initial orbital phase positions as seen in Figure~\\ref{fig:spotseq}. We then computed the expected phase position of each of these 5 anomalies on the 8 following lightcurve residuals assuming that occulted starspots should have moved by $360/8=45\\degr$ in longitude towards the egress (for a prograde and aligned orbit) for every planetary transit. We find that anomalies are always present at all the expected phases.  As the star rotates, the spot surfaces projected on the stellar sphere change; therefore the shapes and amplitudes of these anomalies change between consecutive transits while the same spots are crossed over time. Since a spot surface appears larger at the center of the star compared to the limb, the anomalies due to spots should exhibit larger amplitudes when they are crossed close to the mid-transit time. This is consistent with what is observed in \\koi. For example, anomaly `B' in Figure~\\ref{fig:spotseq} is observed in four consecutive transits. The overall shape of this anolmaly is conserved transit after transit and its amplitude increases before and decreases after the mid-transit.  We conclude that the planet crosses the same spots, transit after transit, at longitudinal positions that differ by 45~\\degr.   These findings imply that the projected spin-orbit angle, $\\lambda$, is very close to 0 for this system. Because occulted spots can be identified transit after transit (Figure~\\ref{fig:spots}), with a change in phase expected from the stellar rotation period, we conclude that the stellar inclination angle is near 90 degrees. This is further supported by the measured $\\vsini=\\koicurSMEvsin$~km/s which is in good agreement with its expected value of $2\\pi\\rstar/P_{\\rm rot} = 4.3$~km/s assuming an inclination angle of 90 degrees. This implies that the true obliquity, i.e. the angle between the stellar rotational axis and a line perpendicular to planet's orbital plane, $\\Psi$, must be close to 0. Finally, we conclude that the planet orbit is prograde since the occulted starspots progress from the transit ingress to the egress limbs.   Assuming that the uncertainty on latitude of a spot corresponds to the planet size, we derive an uncertainty for $\\Psi$ of $\\pm\\arctan(\\rpl/\\rstar)=\\pm10\\degr$. We note that the spot size could be larger than the planet size, but this changes the results only modestly. Assuming a star spot twice the projected size of the planet on the stellar photosphere provides a more conservative uncertainty of $\\pm15\\degr$ for $\\Psi$.  \\cite{winn10} show that stars with transiting planets for which the stellar obliquity is large are preferentially hot (\\teff $> 6250$~K). The low obliquity of \\koib\\ fits this pattern since the stellar temperature is lower than 6250~K.  \\subsubsection{Starspots lifetime} \\label{spotlife}  To help visualize the spot crossings during transit, we subtract the best-fit model from each individual transit and investigate the residuals. We slide a box of duration twice the ingress time across the residuals in each transit epoch, recording the scatter for a given box mid-time (relative to mid-transit) and a given epoch. We then produce an image of this scatter at orbital phase as a function of the epoch number (see Figure~\\ref{fig:spots}). We used an interpolation (with IDL's TRIGRID) with an output sampling of 600 samples in the epochal direction and 300 samples in the transit phase direction. This image reveals individual spots that we define as either ``hot\" or ``cold\" regions, depending on whether the individual slide box residuals are below or above the transit light curve model. The repeating vertical structure is interpreted as spots marching across the transit chord such as seen in the previous section. Each vertical profile is slanted slightly from left to right indicating that the spots progress from the ingress limb to the egress limb. Some spots make their way around the star and reappear again during several stellar rotation periods. For example, the collection of ``cold\" spots in the image starting around Epoch 110 and ending around Epoch 170 seem to be related to the same spot. We conclude from the nearly continuous monitoring of \\koi\\ that the occulted starspots are present on the same stellar chord for at least 100 days, somewhat comparable to the lifetime of sunspots.   As \\kepler\\ continues to monitor transits of hot-Jupiters in front of active stars, it will help to better understand the stellar cycles. If the \\kepler\\ mission is extended, the long term photometry will enable it to produce star-spot maps and learn more about spot mean lifetimes and photospheric differential rotations. In the case of \\koi, we may be able to measure the complete activity cycle for this star and to compare it to another well-know G-dwarf: the Sun.    \\subsubsection{Impact of stellar variability of \\koi\\ on the system parameters}   When the planet transits in front of stellar spots, its transit shape deviates from the averaged phase-folded lightcurve. The effect of occulted stellar spots on the shape of the transit lightcurve is observed in the residuals from the best-fit transit model of the phase-folded light curve (see Figure~\\ref{fig:keplerlc}). Since the stellar activity influences the transit lightcurve profiles, the planetary parameters we derive from these profiles are likely to be affected. This is a well known problem for planets transiting in front of variable stars (e.g. \\citealt{czesla09,desert11a}). Importantly for the present study, the variability affects the stellar density that we assume a fixed value for our determination of the stellar parameters (see Section~\\ref{sec:stellar}). \\cite{czesla09} propose to fit the lower envelope of the transit lightcurve to recover more realistic transit parameters. This assumes that dark stellar structures dominate over bright faculae. In the case of \\koib, we cannot exclude the possibility that every transit is affected by dark or bright stellar regions so that a priori no individual transit light curve can be used as representative of an unaffected profile. Furthermore, because of the stroboscopic effect described above, the phase-folded transit lightcurve possesses combined pattern distortions that prevents the use of its lower envelope to derive more accurate parameters.  We note that the combination of the residuals for the eight transit lightcurves presented in  Figure~\\ref{fig:spotseq} is very similar to the total residuals plotted in Figure~\\ref{fig:keplerlc} and exhibits a symmetrical structure. This symmetrical pattern has a shape which is consistent with what would be expected from an oblate planet but with an amplitude ten to a hundred times larger than expected for a Jupiter size planet \\citep{seager02,barnes03,carter10}. This pattern represents a remarkable coincidence, in that it seems quite symmetric but it is only obtained when all 8 time series are combined, and does not appear in the residuals for any single time series.   \\subsection{Atmospheric constraints for \\koib\\ }\\label{atmo}    The occultation depths measured in each bandpass are combined and turned into an emergent spectrum for the planet. The observed flux of the planet in each bandpass corresponds to the sum of the reflected light and the thermally emitted light. The combined occultation observed in the \\kepler\\ bandpass provides a measure of the planet's geometric albedo. The occultation depth of  $58 \\pm 8$ ppm corresponds to $A_g=0.10 \\pm 0.02$. Assuming a Lambertian criterion $A_{\\rm B} \\leq 1.5 \\times A_{\\rm g}$, and assuming that the geometric albedo measured in the \\kepler\\ bandpass is the same at every wavelength, we infer an upper limit to the Bond albedo of $A_{\\rm B} \\leq 0.18$ at the 1-$\\sigma$ level. This result is conservative in the sense that the true albedo may be even lower, if some of the occultation signal we have measured in the \\kepler\\ bandpass is due to thermal emission rather than reflected light. This is expected for the temperature regime of hot-Jupiters. Therefore, the Bond albedo could be well below 0.18.   We estimate the thermal component of the planet's emission in the two \\spitzer\\ band-passes from the occultation depths measured at 3.6 and 4.5~\\micron\\ (see Table~\\ref{tab:spitzer}). We assume that the planetary emission is well reproduced by a black-body spectrum and translate the measured depth of the secondary eclipse into brightness temperatures. We use the PHOENIX atmospheric code \\citep{hauschildt99} to produce theoretical stellar models for the star \\koi. Taking the \\spitzer\\ spectral response function into account, the ratio of areas of the star and the planet and the stellar spectra, we derive the brightness temperatures that best-fit the observed eclipse depths measured in the two IRAC bandpasses. The brightness temperature calculated this way result in T$_{\\rm{3.6~\\micron}}$=$1880\\pm110$~K and T$_{\\rm{4.5~\\micron}}$=$1770\\pm150$~K.   We compare our data to the hot-Jupiter atmospheric model described in \\cite{fortney08} which has been used for a variety of close-in planets \\citep{fortney05,fortney06}. We aim at broadly distinguishing between different classes of model atmospheres, given that there are no prior constraints on the basic composition and structure. The atmospheric spectrum calculations are performed for 1D atmospheric pressure-temperature (P-T) profiles and use the equilibrium chemical abundances, at solar metallicity, described in \\cite{lodders02, lodders06}. This is a self-consistent treatment of radiative transfer and chemical equilibrium of neutral species. The opacity database is described in \\cite{freedman08}. The models are calculated for various values of the dayside-to-nightside energy redistribution parameter ($f$) and allow for the presence of TiO at high altitude, which may play the role of an absorber and which likely lead to an inversion of the P-T profile.   We compare the data to model predictions and select models with the best reduced \\chisq. We note that this is not a fit involving adjustable parameters. All models shown assume redistribution of absorbed flux over the dayside only, favoring low values for $f$. Since our observations are only weakly constraining, the 3.6 and 4.5 $\\mu$m depth ratios are best fitted by a model without a temperature inversion, although the model with a temperature inversion is a reasonable fit as well. In the model with no inversion the optical opacity sources are neutral atomic Na and K, while in the inverted model the optical opacity is dominated by gaseous TiO and VO molecules. The relatively large occultation depth in the Kepler bandpass is not matched by these simple models.  This may indicate that silicate clouds (as is seen in L-type brown dwarfs, but not modeled here) may lead to a higher amount of scattered flux from the planet. It may also indicate that optical opacity sources are overestimated, such that more planetary thermal flux is being seen from deeper layers. The growing sample size of Kepler occultation detections may show trends with planetary \\teff\\ that could help to clarify this issue.  \\cite{hartman10} shows that there is a correlation between the surface gravity of hot Jupiters and the activity levels of the host stars, such that high surface gravity planets tend to be found around high-activity stars. The position of \\koib\\ on such a diagram is consistent with this trend. \\cite{knutson10} show that there could also be a correlation between the host star activity level and the thermal inversion of the planetary atmosphere. In this picture, the strong XEUV irradiation from the active stellar host of a hot-Jupiter depletes the atmosphere of chemical species responsible for producing inversions. The \\logrhk\\ of \\koi\\ is consistent with the case of no thermal inversion in the framework developed by \\cite{knutson10}. We also evaluated the empirical index defined by \\citet{knutson10}, which could be correlated with the presence of a thermal inversion. Using the same definition, we find indices of 0.026\\% above and 0.034\\% below the predicted black-body fluxes in the 3.6 and 4.5 micron bands. These indices suggest that the atmosphere of \\koib\\ is consistent with a non-inverted profile, which is in agreement with the results of our present study."
    },
    "1107/1107.4610_arXiv.txt": {
        "abstract": "The thermodynamic and dynamical properties of a variable dark energy model with density scaling as $\\rho_{x} \\propto (1+z)^{m}$, z being the redshift, are discussed following the outline of Jetzer et al. \\cite{Jetzer+11}. This kind of models are proven to lead to the creation/disruption of matter and radiation, which affect the cosmic evolution of both matter and radiation components in the Universe. In particular, we have concentrated on the temperature-redshift relation of radiation, which has been constrained using a very recent collection of cosmic microwave background (CMB) temperature measurements up to $z \\sim 3$. For the first time, we have combined this observational probe with a set of independent measurements (Supernovae Ia distance moduli, CMB anisotropy, large-scale structure and observational data for the Hubble parameter), which are commonly adopted to constrain dark energy models. We find that, within the uncertainties, the model is indistinguishable from a cosmological constant which does not exchange any particles with other components. Anyway, while temperature measurements and Supernovae Ia tend to predict slightly decaying models, the contrary happens if CMB data are included. Future observations, in particular measurements of CMB temperature at large redshift, will allow to give firmer bounds on the effective equation of state parameter \\weff\\ of this kind of dark energy models. ",
        "introduction": "The current standard cosmology model relies on the existence of two unknown dark components, the so called ``dark matter'' (DM) and ``dark energy'' (DE), which amount to $\\sim 25\\%$ and $\\sim 70\\%$ of the total energy budget in the Universe, respectively. According to several observations, the Universe is spatially flat and in an accelerated phase of its expansion \\cite{perl, reiss, debernardis, spe07, caldwell}. DE, described as a cosmological constant $\\Lambda$ in its simplest form, is modelled by a fluid with a negative pressure, which is a fundamental ingredient to explain the actual accelerated expansion of the Universe.   Several models have been proposed to explain DE \\cite{PR03, Pad03, D+05, CTTC06, Caldwell02, PR88, RP88, SS00}. An alternative consists to consider a phenomenological variable DE density with continuous creation/disruption of photons \\cite{lima+96, LA99, lima+00, puy, Jetzer+11} or matter \\cite{CW90, ma}. The DE might decay/grow slowly in the course of the cosmic evolution and thus provide the source/sink term for matter and radiation. Different such models have been discussed and strong constraints come from very accurate measurements of the cosmic microwave background (CMB) radiation and other typical cosmological probe.   CMB radiation is the best evidence for an expanding Universe starting from an initial high density state. Within the Friedmann-Robertson-Walker (FRW) models of the Universe the radiation, after decoupling, expands adiabatically and scales as $(1+z)$, $z$ being the redshift \\cite{Weinberg72}. If we assume that each component is not conserved, contrarily to the standard scenario, then depending on the decay mechanism of the DE, the created photons could lead to distortions in the Planck spectrum of the CMB, and change the evolution of its temperature. The chance to appreciate the deviation from the standard temperature evolution is given by the increasingly number of recent works collecting observations of CMB temperature both at low \\cite{battistelli, luz09} and higher redshifts \\cite{Noterdaeme+11}.  Following the theoretical lines of \\cite{lima+96, LA99, lima+00, puy}, in Jetzer et al. \\cite{Jetzer+11} we have discussed a variable DE model $\\Lambda (z) \\propto (1+z)^{m}$ decaying into photons and DM particles. In particular we studied thermodynamical aspects in the case of a continuous photon creation, which implies a modified temperature redshift relation for the CMB. We have tested the predicted temperature evolution of radiation with some recent data on the CMB at higher redshift from both Sunyaev-Zel'dovich (SZ) effect and high-redshift QSO absorption lines.  In this paper we will present the results obtained in \\cite{Jetzer+11} within a more detailed theoretical background, by further discussing the main hypothesis which concern the energy conservation and thermodynamical laws. Furthermore, we will test our model with an updated collection of data, joining our sample with the new temperature observations reported in \\cite{Noterdaeme+11}. For the first time we will combine to the temperature measurements also the data from different kind of observations, like distance moduli of Supernovae Ia, observations of the CMB anisotropy and the large-scale structure, together with observational Hubble parameter estimations. We will dedicate a particular attention to the estimate of both the actual matter density parameter, the DE parameter $m$ and the effective equation of state parameter $w_{\\rm eff}$. The impact of each observational probe on these estimates is also investigated.  The paper is organised as follows. In \\S \\ref{sec:theory} we will present the model, while the data sets and fitting procedure are described in \\S \\ref{sec:data}. The results are discussed in \\S \\ref{sec:results} and \\S \\ref{sec:conclusions} is devoted to a discussion of the results and conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have presented the properties of a variable DE model, following the approach in Jetzer et al. \\cite{Jetzer+11} (see also \\cite{lima+96, LA99, lima+00, puy, ma}). The model relies on the assumption that the DE component is not conserved and exchanges particles with both DM and radiation components. In particular, we show that in order to have a viable model, DE needs to decay besides into radiation also in matter. Motivated by these considerations, we have concentrated the analysis on the case when only a small fraction is exchanged with radiation. We have dedicated particular attention to the thermodynamic properties of the model, discussing the theoretical relation between the radiation temperature and redshift and we have matched it to the most updated collection of CMB temperature measurements \\cite{luz09, Jetzer+11, Noterdaeme+11}, consisting of both SZ data and high-redshift QSO adsorption line observations. First, we have constrained the model by setting $\\gamma = 4/3$ and finding $m = 0.03^{+0.08}_{-0.09}$, consistent with our previous estimate within uncertainties, still consistent with the standard case of a $T \\propto (1+z)$, i.e. with an effective equation of state $\\weff = -1$. When $\\gamma$ is left free, then $m= 0.25_{-0.17}^{+0.23}$. Although the large contours indicate that is not possible to find statistically relevant departures from a standard cosmological constant, we find that a decaying DE, with an effective equation of state $\\weff > -1$ is preferred. Then, for the first time, we test this kind of model combining both CMB temperature measurements and different observational data sets, like Supernovae Ia, CMB and large-scale structure data, etc. We find that \\Om\\ is almost independent on the combination of data we use, and in the best case is constrained to $\\Om = 0.27_{-0.01}^{+0.02}$. Moreover, while temperature measurements and SN data tend to furnish $m \\gsim 0$, consistently with a decaying DE model, other datasets, like CMB data prefer the opposite situation with $m \\lsim 0$, i.e. a phantom-like model with $\\weff < -1$.  Although the present data do not allow to find strong discrepancies with a classical cosmological constant model, we think that future surveys at high redshift could collect further measurements for CMB temperature, which could help us to further constrain the temperature-redshift relation. This kind of information, together with larger and higher quality samples of Supernovae and better CMB data or other standard candles like Gamma ray burst (GRBs) can allow to give independent constraints on both matter density \\Om\\ and effective equation of state \\weff.  We notice that if DE does not decay into radiation (corresponding to $\\epsilon=\\tilde{\\epsilon}=0$ and thus $m$ is no longer constrained) then the CMB temperature will scale in the standard way. Clearly, this would imply that if DE decays, this has to be into DM only. On the other hand a deviation  of the CMB temperature from the standard scaling could be interpreted as DE decaying also into radiation. In which case with Eq. (\\ref{ma3}) one can determine $m$ and/or $\\tilde{\\epsilon}$ and thus get some insights on the decay mode of DE into radiation. Future data on the CMB temperature will allow to shed light on this important issue."
    },
    "1107/1107.5117.txt": {
        "abstract": "We conducted very long baseline interferometry~(VLBI) observations of five radio-loud narrow-line Seyfert~1~(NLS1) galaxies in milliarcsecond resolutions at 1.7~GHz~($\\lambda$18~cm) using the Very Long Baseline Array~(VLBA).  Significant parsec-scale structures were revealed for three out of the five sources with high brightness temperature by direct imaging; this is convincing evidence for nonthermal jets.  FBQS~J1644+2619 with an inverted spectrum showed a prominent one-sided linear structure, indicating Doppler beaming with an intrinsic jet speed of $>0.74c$.  FBQS~J1629+4007, also with an inverted spectrum, showed rapid flux variability, indicating Doppler beaming with an intrinsic jet speed of $>0.88c$.  Thus, we found convincing evidence that these two NLS1s can generate at least mildly or highly relativistic jets, which may make them apparently radio loud even if they are intrinsically radio quiet.  On the other hand, the other three NLS1s had steep spectra and two of them showed significantly diffuse pc-scale structures, which were unlikely to be strongly beamed.  Thus, some NLS1s have the ability to generate jets strong enough to make them intrinsically radio loud without Doppler beaming.  NLS1s as a class show a number of extreme properties and radio-loud ones are very rare.  We build on these radio results to understand that the central engines of radio-loud NLS1s are essentially the same as that of other radio-loud AGNs in terms of the formation of nonthermal jets. ",
        "introduction": "\\label{section:introduction} Active galactic nuclei~(AGNs) are believed to be powered by gravitational energy released from accreting matter onto supermassive black holes~(SMBHs) at the center of galaxies.  Narrow-line Seyfert~1 galaxies~(NLS1s) are a subclass of AGNs and are defined by the following optical spectral properties: (1)~permitted lines are slightly broader than forbidden lines; (2)~the full width at half-maximum~(FWHM) of the H$\\beta$ emission line is $<$2000~km~s$^{-1}$; and (3)~the flux ratio [\\ion{O}{3}]/H$\\beta<$3, but exceptions are allowed when strong high-ionization iron lines such as [\\ion{Fe}{7}] $\\lambda$6087 and [\\ion{Fe}{10}] $\\lambda$6375, which are seldom seen in type-2 AGNs, are present~\\citep{Osterbrock&Pogge1985,Pogge2000}. Other extreme properties of NLS1s are strong permitted Fe $_\\mathrm{II}$ emission lines in optical and ultraviolet spectra \\citep{Boroson&Green1992}, rapid and large-amplitude X-ray variability \\citep{Leighly1999a}, a steep soft X-ray spectrum \\citep{Boller_etal.1996,Brandt_etal.1997,Grupe_etal.1998,Grupe_etal.2004}, and a frequently observed blueshifted line profile \\citep{Zamanov_etal.2002,Boroson2005,Leighly&Moore2004}.   In terms of radio luminosity ($10^{20}$--$10^{23}$~W~Hz$^{-1}$) and the compact source size ($\\lesssim300$~pc; \\citealt{Ulvestad_etal.1995,Moran2000}), the radio continuum properties of NLS1s and normal Seyfert~1 galaxies are not remarkably different.  NLS1s are generally radio-quiet objects\\footnote{Radio loudness, $R$, is conventionally defined as the ratio of the 5-GHz radio to {\\it B}-band flux densities, with a threshold of $R=10$ separating the radio-loud and radio-quiet objects \\citep{Visnovsky_etal.1992,Stocke_etal.1992,Kellermann_etal.1994}.}.  The fraction of radio-loud objects in NLS1s is low \\citep{Stepanian_etal.2003,Greene_etal.2006,Zhou_etal.2006}, that is, $\\sim7$\\% ($R>10$) and $\\sim2.5$\\% ($R>100$) (\\citealt{Komossa_etal.2006a}, see also \\citealt{Zhou&Wang2002,Yuan_etal.2008}), which is significantly lower than normal Seyfert galaxies and quasars \\citep[$\\sim10$\\%--15\\%;][]{Ivezic_etal.2002}.   There is increasing evidence that the unusual optical/UV/X-ray properties of NLS1s are physically related to the high mass accretion rates close to the Eddington limit \\citep{Boroson&Green1992,Pounds_etal.1995,Sulentic_etal.2000,Boroson2002} of the relatively low-mass black holes~($\\sim 10^5$--10$^7\\ M_\\sun$), which are located at the lower-mass end of the known population of SMBHs \\citep[][and references therein]{Peterson_etal.2000,Grupe&Mathur2004,Komossa&Xu2007}.  Hence, NLS1s are expected to be promising samples for understanding the nature of the extreme accretion onto rapidly growing SMBHs.  The radio emission of AGNs and X-ray binaries is suggested to be also related to the black hole mass and accretion rate \\citep{Merloni_etal.2003,Heinz&Sunyaev2003,Falcke_etal.2004}.  Many studies using large samples show a correlation between radio loudness and the black hole mass \\citep{Laor2000,Lacy_etal.2001,Dunlop_etal.2003,McLure&Jarvis2004,Metcalf&Magliocchetti2006,Woo&Urry2002} and an anti-correlation between radio loudness and the accretion rate \\citep{Ho2002,Lacy_etal.2001,Maccarone_etal.2003,Greene_etal.2006}.  These correlations predict that AGNs with high accretion rates and low-mass black holes such as NLS1s should be radio quiet.  In fact, as mentioned above, most NLS1s have been detected as radio-quiet objects.  In the analogy of X-ray binaries, NLS1s are at the high/soft state or at the very high state \\citep{Mineshige_etal.2000,Wang&Netzer2003}, where radio emission is quenched \\citep{Maccarone_etal.2003}.  Radio-loud NLS1s are rare, but they do exist \\citep{Siebert_etal.1999,Grupe_etal.2000,Zhou&Wang2002,Zhou_etal.2003,Whalen_etal.2006,Komossa_etal.2006a,Komossa_etal.2006b,Yuan_etal.2008}.  The radio powers of the radio-loud NLS1s exceed those of the most luminous starburst galaxies, suggesting that strong AGN jet emissions are expected to be responsible for making these NLS1s radio loud.  Although the black hole masses of the radio-loud NLS1s ($\\sim 10^7\\ M_\\sun$) may be slightly larger than those of the radio-quiet NLS1s, the radio-loud NLS1s are still at a lower mass regime in AGNs and at a very high accretion rate regime (\\citealt{Komossa_etal.2006a}, see also \\citealt{Whalen_etal.2006,Yuan_etal.2008}); although such properties tend toward radio-quietness.  One possible explanation of the existence of radio-loud NLS1s is that the Doppler beaming effect on nonthermal jets influences radio loudness as well as some other radio-loud AGNs.  However, most radio-loud NLS1s show steep soft-excess X-ray spectra, which are the same as those of the radio-quiet NLS1s \\citep{Komossa_etal.2006a,Yuan_etal.2008} and in contrast to the flat X-ray spectra due to the inverse Compton process in relativistic jets in radio-loud quasars~\\citep{Reeves_etal.1997}.  The steep X-ray spectra of radio-loud NLS1s may be synchrotron tails originating from the relativistic jets \\citep{Zhou_etal.2007,Yuan_etal.2008}, as seen in high-energy-peaked BL~Lac objects~\\citep{Padovani&Giommi1995}.  On the other hand, \\citet{Komossa_etal.2006a} indicated that many radio-loud NLS1s show steep radio spectra where strong beaming is generally not expected, although {\\it very} radio-loud NLS1s may systematically show flat radio spectra \\citep{Yuan_etal.2008}.  Several radio-loud NLS1s show physically unrealistic, high brightness temperatures, derived from radio flux variability.  This indicates the Doppler beaming effect if not interstellar scattering \\citep{Zhou_etal.2003,Yuan_etal.2008}.  However, the Doppler beaming effect on the jets of NLS1s is not yet well established.  Very long baseline interferometry~(VLBI) is one of the most powerful tools for revealing the radio properties in parsec~(pc) scales by direct imaging at milliarcsecond~(mas) resolutions.  Arcsecond-resolution observations have resolved the structures of only a few NLS1s because they are generally quite compact radio sources \\citep{Ulvestad_etal.1995,Moran2000}.  VLBI observations on a large number of radio-quiet and radio-loud NLS1s are crucial for understanding the nature of the jet phenomena in these central engines.  VLBI images of only several objects have been reported; that is, for radio-quiet NLS1s: NGC~4395 \\citep{Wrobel_etal.2001}, MRK~766, AKN~564 \\citep{Lal_etal.2004}, and NGC~5506 \\citep{Middelberg_etal.2004} and for radio-loud NLS1s: \\object{SDSS J094857.31+002225.4} (\\citealt{Doi_etal.2006a}, hereafter \\citetalias{Doi_etal.2006a}), J16290+4007, J16333+4718, and B3~1702+457 \\citep{Gu&Chen2010}.  SDSS~J094857.31+002225.4, which is the most radio-loud narrow-line object with $R\\approx2000$, is suggested to be Doppler beamed by radio flux variability \\citep{Zhou_etal.2003}, and it has been revealed that Doppler-boosted jets are required to explain the directly measured high brightness temperatures of its very compact radio emissions seen in the VLBI images \\citepalias{Doi_etal.2006a}.  The recent discovery of $\\gamma$-ray emission from SDSS~J094857.31+002225.4 by the {\\it Fermi Gamma-Ray Space Telescope} \\citep{Abdo_etal.2009a} and the subsequent VLBI observations have established the presence of a relativistic jet in this extreme narrow-line quasar \\citep{Abdo_etal.2009b,Giroletti_etal.2011}.   We conducted VLBI observations for five radio-loud NLS1s at 1.7 and 8.4~GHz using the Very Long Baseline Array~(VLBA) and the Japanese VLBI Network~(JVN), respectively.  The results of the JVN observations, which have been published (\\citealt{Doi_etal.2007}, hereafter \\citetalias{Doi_etal.2007}), showed no significant radio structure in these five objects.  In the present paper, we report the observations at 1.7~GHz, where synchrotron radio jet structures are generally expected to be more prominent than at higher frequencies.  Based on the VLBA observations at 5~GHz for three radio-loud NLS1s that are three of our five samples~(for one source also at 2.3 and 8.4 GHz), \\citet{Gu&Chen2010} reported that compact structures predominate radio emissions.  The present paper is structured as follows.  The sample selection is shown in Section~\\ref{section:sample}.  Section~\\ref{section:observation} describes our observations and the data reduction procedures.  In Section~\\ref{section:result}, we present the results and image analyses, and in Section~\\ref{section:discussion}, we discuss the observed pc-scale jet properties and their implications.  Finally, the summary is presented in Section~\\ref{section:summary}.  Throughout this paper, a flat cosmology is assumed, with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\mathrm{M}=0.27$, and $\\Omega_\\mathrm{\\Lambda}=0.73$ \\citep{Spergel_etal.2003}. ",
        "conclusions": "\\label{section:discussion} We found significantly resolved pc-scale radio structures for three NLS1s and compact structures for two NLS1s.  The key results from our VLBA images are as follows:  (1)~detections of very high brightness temperatures and linear structures of the pc-scale radio emissions, which provide evidence of nonthermal jets in the NLS1s~(Section~\\ref{section:nonthermaljets}); (2)~asymmetry of radio structures and core spectral indices that are very important information for evaluating the Doppler beaming effect on the jets (Section~\\ref{section:asymmetricstructure}); and (3)~discovery of different radio properties of the five sources, which suggest why these NLS1s are radio loud~(Section~\\ref{section:naturesofRLNLS1s}).  Detailed descriptions of individual objects are given in the Appendix.    \\subsection{Strong Evidence for Nonthermal Jets}\\label{section:nonthermaljets} The previous JVN observations at 8.4~GHz detected brightness temperatures of at least $10^{7.4}$~K and high radio luminosities, which cannot be attributed to any stellar process; radio emissions presumably originate from nonthermal jets.  The VLBA observations have also revealed that radio components with very high brightness temperatures, which provide stronger constraints, i.e., $T_\\mathrm{B}>10^{8.7}$~K, were associated with the five NLS1s.  Significant pc-scale radio structures from the VLBA observations were newly detected, providing stronger evidence for the presence of nonthermal jets.  \\object{FBQS J1644+2619} and B3~1702+457 clearly displayed linear structures. (RX~J0806.6+7248 also showed a linear structure in the east and west of its nuclear region.)  RX~J0806.6+7248, \\object{RX J1633.3+4718}, and \\object{B3 1702+457} are steep-spectrum sources~($\\alpha < -0.5$) with very high brightness temperatures.  This is a strong evidence for a nonthermal process dominating the radio emissions.  On the other hand, \\object{FBQS J1629+4007} and FBQS~J1644+2619 showed inverted spectra~($\\alpha>0$).  The observed brightness temperatures were too high for free--free emissions, although a radio spectrum of not only nonthermal synchrotron emissions but also thermal free--free emissions could be inverted in an optically thick regime\\footnote{For example, a thermal torus of at most $T_\\mathrm{B} \\sim 10^6$~K such as the one inferred for type-2 Seyfert NGC~1068 \\citep{Gallimore_etal.1997,Gallimore_etal.2004} was imaged using the VLBA.}.  We consider a putative source with having a diameter of lower than that of the VLBA beams radiating the free--free emissions at electron temperatures equal to the observed brightness temperatures and being optically thick at our observing frequency.  Under these conditions, the X-ray luminosities at 0.1--2.4~keV of the bremsstrahlung emission would be $\\sim10^{48}$~erg~s$^{-1}$, which cannot meet the observed X-ray luminosities $\\sim10^{44}$~erg~s$^{-1}$ (Table~\\ref{table1}); these were detected in the {\\it ROSAT} All Sky Survey \\citep{Bade_etal.1995,Laurent-Muehleisen_etal.1998} and should include the contribution not only of bremsstrahlung emission but also the inverse Compton and synchrotron emissions.  Hence, for all observed NLS1s, we found strong evidence that the nonthermal process dominates the observed radio emissions.  Because the VLBI-correlated fluxes at 1.7~GHz were equivalent to much of the total fluxes in the VLA at 1.4~GHz, we conclude that the radio emissions of all these radio-loud NLS1s are dominated by the nonthermal synchrotron process in jets and that these pc-scale jets make the NLS1s radio loud.    \\subsection{Asymmetric Core--Jet Structures}\\label{section:asymmetricstructure} The AGN radio sources generally consist of steep-spectrum jets ($\\alpha<0$) and a flat-/inverted-spectrum core ($\\alpha \\ga 0$); the core is the location where the synchrotron emission becomes visible from an optically thick upstream to an optically thin downstream at an observing frequency.  The core flux contribution generally tends to a relative increase in the higher frequency radio images; therefore, the JVN emissions should originate from cores.  A structural asymmetry relative to the core was found in the limited dynamic ranges of our images, although only two components were seen in FBQS~J1629+4007 and RX~J1633.3+4718.  The possible causes of the asymmetry are as follows: (1)~the jets are intrinsically one-sided or more lossy (and hence radiative) on one side than the other; or (2)~Doppler boosting effect on the relativistic jets; or (3)~the jets are intrinsically symmetric but appear asymmetric owing to the stronger free-free absorption toward the counter-jets due to an obscuring ionized gas intervening along our line of sight \\citep{Walker_etal.2000,Jones et al.2001,Kameno_etal.2001}.  Both (2) and (3) require that the stronger side should be approaching jets, while the weaker side should be receding jets.  In this paper, we discuss the possible parameters of the jets in the NLS1s on the assumption that the jet are intrinsically symmetric.  The images of RX~J0806.6+7248 and B3~1702+457 showed possible counter features, and thus indicating receding jets.  However, it should be noted that in some cases, the VLBI imaging processes probably generate a mimic counter component near the core \\citep[e.g., imaging demonstrations by][]{Carilli_etal.1994}.  We cannot be confident that the counter-jet features were genuine receding jets (the counter jets in RX~J0806.6+7248 may exist because a diffuse counter feature, which is misaligned with the linear structure, was also seen at a small distance; the Appendix.  On the other hand, no counter component can be seen in FBQS~J1644+2619, which has a significant linear jet structure in the limited image dynamic range; this result provides a reliable constraint for a receding jet.  Hereafter, we consider FBQS~J1644+2619 for discussing asymmetry.       \\subsubsection{Doppler Beaming Effect Estimated from One-Sidedness}\\label{section:Dopplerbeamingeffect} In case Doppler boosting was responsible for the observed asymmetry of the jet structure, we can limit the jet speed and inclination by comparing the approaching and receding jets on the assumption that the rest-frame radio spectra and intrinsic jet speeds of both the sides were the same.  A bulk Lorentz factor is defined as $\\Gamma \\equiv (1-\\beta^2)^{-0.5}$, where $\\beta \\equiv v/c$ ($v$ and $c$ are the speeds of the source and light, respectively).  The Doppler factor is defined as \\begin{equation} \\delta \\equiv \\frac{1}{\\Gamma (1-\\beta\\cos{\\Phi})}, \\end{equation} where $\\Phi$ is the jet viewing angle, which is the angle between the direction of the source velocity and our line of sight \\citep[][Appendix~B]{Ghisellini_etal.1993}.  The frequency and the flux density at the rest frame, $\\nu_\\mathrm{rest}$ and $S_{\\nu\\mathrm{rest}}^\\mathrm{rest}$, are observed as those at the observer frame, $\\nu_\\mathrm{obs} = \\delta \\nu_\\mathrm{rest}$ and $S_{\\nu\\mathrm{obs}}^\\mathrm{obs} = \\delta^3 S_{\\nu\\mathrm{rest}}^\\mathrm{rest}$, respectively.  Assuming the spectral index $\\alpha$, $S_{\\nu\\mathrm{obs}}^\\mathrm{obs} = \\delta^{3-\\alpha} S_{\\nu\\mathrm{obs}}^\\mathrm{rest}$ \\citep{Rybicki&Lightman1979}.  Hence, the flux-density ratio of the approaching and receding jets is \\begin{eqnarray} r &=& \\frac{S_{\\nu\\mathrm{obs}}^\\mathrm{rest}(\\beta, \\Phi)}{S_{\\nu\\mathrm{obs}}^\\mathrm{rest}(\\beta, \\Phi + \\pi)}\\nonumber\\\\ &=& \\left( \\frac{1+\\beta \\cos{\\Phi}}{1-\\beta \\cos{\\Phi}} \\right)^q, \\end{eqnarray} where $q=3-\\alpha$; when considering the brightness ratio instead of the flux-density ratio in continuous jets $q=2-\\alpha$ is applied \\citep{Ghisellini_etal.1993}.  Therefore, we obtained \\begin{equation} \\beta \\cos{\\Phi} > \\frac{r^{1/q}-1}{r^{1/q}+1}.\\label{eq:fluxratio} \\end{equation} Using this inequality, we can estimate the lower limit of $\\beta$ and the upper limit of $\\Phi$.    \\begin{table} \\caption{Parameters of the Doppler Beaming Effect for FBQS~J1644+2619\\label{table4}} \\begin{center} \\begin{tabular}{ccccccc} \\hline\\hline Source name & Comparison & $r$ & $k$ & $\\alpha$ & $\\beta_\\mathrm{min}$ & $\\Phi_\\mathrm{max}$ \\\\ &  &  &  &  &  & ($\\degr$) \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) \\\\\\hline \\multicolumn{1}{l}{FBQS~J1644+2619} & $S_{\\nu}(\\mathrm{core})/3\\sigma$  & $>$147 & 3  & $\\alpha_I$ & 0.74  & 42  \\\\ $\\uparrow$ & $I_{\\nu}(\\mathrm{jet})/3\\sigma$  & $>$37 & 2  & $\\alpha_S$ & 0.77  & 39  \\\\\\hline \\end{tabular} \\end{center} \\tablecomments{Col.~(1)~source name; Col.(2)~measurements compared to derive the ratio of jet to counter jet.  $S_{\\nu}(\\mathrm{core})$ means the core's flux density.  $3\\sigma$ means three times the rms~image noise as an upper limit for a counter jet.  $I_{\\nu}(\\mathrm{jet})$ means the intensity of the jet emission at the position separated from the peak by the spatial resolution along the jets; Col.(3)~ratio of jet to counter jet; Col.~(4)~$k$ value used for the calculation of the Doppler beaming effect, where $q \\equiv k- \\alpha$ (Section~\\ref{section:Dopplerbeamingeffect}); Col.~(5)~assumed spectral index (Table~\\ref{table3}); Col.~(6) lower limit of jet speed (Equation~(\\ref{eq:fluxratio})); Col.~(7)~upper limit of the angle of jet axis (Equation~(\\ref{eq:fluxratio})).} \\end{table}  \\begin{table*} \\caption{Parameters of the Variability Doppler-beaming Effect for FBQS~J1629+4007\\label{table5}} % {\\tiny \\begin{center} \\begin{tabular}{ccccccccccccc} \\hline\\hline Name & $I^\\mathrm{FIRST}_\\mathrm{1.4}$ & Obs.~Date & $S^\\mathrm{NVSS}_\\mathrm{1.4}$ & Obs.~Date & Variability & $\\log{T_\\mathrm{B}}$ & $\\log{T_\\mathrm{int}^\\mathrm{max}}$ & $\\delta_\\mathrm{min}$ & $\\Phi_\\mathrm{max}$ & $\\beta_\\mathrm{min}$ & $\\log{P_\\mathrm{1.7GHz}^\\mathrm{int}}$ & $\\log{R_\\mathrm{*}^\\mathrm{int}}$ \\\\ & (mJy beam$^{-1}$) &  & (mJy) &  &  & (K) & (K) &  & ($\\degr$) &  & (W Hz$^{-1}$) &  \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11) & (12) & (13) \\\\\\hline \\multicolumn{1}{l}{FBQS~J1629+4007} & 11.97 $\\pm$ 0.17 & 1994 Aug 21 & 9.0 $\\pm$ 0.5 & 1995 Apr 16 & $5.6\\sigma$ & 13.0  & 11.2  & 4.0  & 14  & 0.88  & $<$22.7 & $<$0.74 \\\\ $\\uparrow$ & $\\uparrow$ & $\\uparrow$ & $\\uparrow$ & $\\uparrow$ & $\\uparrow$ & $\\uparrow$ & 12 & 2.1  & 28  & 0.63  & $<$23.4 & $<$1.43 \\\\\\hline \\end{tabular} \\end{center} % } \\tablecomments{Col.~(1)~source name; Col.(2)~1.4-GHz peak flux density from FIRST; Col.~(3)~observation date of the FIRST data; Col.~(4)~1.4-GHz peak flux density from NVSS; Col.~(5)~observation date of the NVSS data; Col.~(6)~significance of the detection of variability; Col.~(7)~brightness temperature derived from variability; Col.~(8)~assumption of maximum intrinsic brightness temperature (Section~\\ref{section:fluxvariability}); Col.~(9)~lower limit of Doppler factor; Col.~(10)~upper limit of viewing angle of jet axis; Col.~(11)~lower limit of jet speed; Col.~(12)~intrinsic radio power; Col.~(13)~intrinsic radio loudness.} \\end{table*}    Using the inequality, we derived the parameters of the Doppler effect for FBQS~J1644+2619 (Table~\\ref{table4}) without considering the free--free absorption.  We considered that the core was the first visible component in the approaching jet stream and that the flux density of the counter component in the receding stream was less than three times the rms~image noise (Table~\\ref{table2}).  We found that $\\beta>0.74$, which is either a mildly or highly relativistic jet speed and that $\\Phi<42$\\arcdeg, which is a mild constraint for the inclination in this NLS1 (Equation~(\\ref{eq:fluxratio})).  For the case of the brightness ratio in continuous jets ($q=2-\\alpha$), we found a similar result of $\\beta>0.77$ and $\\Phi<39$\\arcdeg.  This result indicates that the central engine of an NLS1 can generate nonthermal jets with quite a high speed.  The radio loudness of this NLS1 is potentially due to the Doppler beaming effect.  We cannot obtain $\\delta$ by only considering the apparent jet structure.  VLBI image monitoring will provide the apparent jet speed in order to solve both $\\beta$ and $\\Phi$ for obtaining $\\delta$.  Another method to obtain the constraint $\\delta$ by imaging is to confirm the compactness of the source at a higher spatial resolution.  For example, by VLBI observations, the very radio-loud ($R \\approx 2000$) NLS1 SDSS~J094857.31+002225.4 has highly relativistic jets viewed pole-on because of its compactness, which is indicative of an extremely high brightness temperature requiring Doppler boosting of $\\delta>$2.7--5.5 \\citepalias{Doi_etal.2006a}.     \\subsubsection{Free--Free Absorption}\\label{section:freefreeabsorption} The possibility of the apparent jet asymmetry because of the free-free absorption due to the intervening ionized thermal plasma is discussed for FBQS~J1644+2619.  The optical depth of free--free absorption is \\begin{equation} \\tau_\\mathrm{ff} \\approx 8.235 \\times 10^{-2} T_\\mathrm{e}^{-1.35} \\nu^{-2.1} \\int n_\\mathrm{e} dl, \\end{equation} where $T_\\mathrm{e}$~(K) and $n_\\mathrm{e}$~(cm$^{-3}$) are the electron temperature and density of the absorbing gas, respectively, $\\nu$ is the frequency~(GHz), and $l$ is the path length in pc~\\citep{Mezger_Henderson1967}.  The attenuation of more than a factor of 37 (Table~\\ref{table4}) that corresponds to $\\tau_\\mathrm{ff}>3.6$ is required to make a putative counter-jet opaque.  As a rough estimate, considering a path length equivalent to the length of the putative counter-jet of 20~mas~(50~pc), the fully ionized gas with $T_\\mathrm{e}>8000$~K, and $\\nu=(1+z)\\times1.7$~GHz, a large amount of electrons of the column density of electrons $N_\\mathrm{e}(=n_\\mathrm{e}l) > 10^{23}$~cm$^{-2}$ would be required in the nuclear region.  {\\it Chandra} X-ray observations of FBQS~J1644+2619 showed that the fitted absorption column density $N_\\mathrm{H}$ (natural gas) is almost identical to the Galactic value $5.12 \\times 10^{20}$~cm$^{-2}$ \\citep{Yuan_etal.2008}, suggesting a very poor intervening medium toward the nucleus.  Thus, such a large free--free absorber is unlikely in the FBQS~J1644+2619 nucleus unless an extremely ionized fraction is assumed.  Furthermore, the VLA image of FBQS~J1644+2619 (Figure~\\ref{figure1}(d)) also seems to show an asymmetric structure, which is presumably related to the AGN activity~(the Appendix).  The large-scale asymmetry up to 30~kpc implies that the free-free absorption cannot be responsible for the apparent asymmetry, because such optically thick plasma would provide thermal radio emission of $\\sim4$~Jy (Equation~(\\ref{equation:brightnesstemperature})), which is much larger than the observed total flux density.  Hence, we conclude that Doppler beaming should be preferentially efficient and it contributes to making FBQS~J1644+2619 one-sided, unless the jets are intrinsically one-sided.  The intriguing large structure of FBQS~J1644+2619 is discussed later in Section~\\ref{section:naturesofRLNLS1s}.     \\subsection{Doppler Beaming Inferred from Flux Variability}\\label{section:fluxvariability} The Doppler factor can also be restricted by brightness temperature, which is derived from flux variability.  For our sample, we checked the flux variation between the NRAO VLA Sky Survey~(NVSS; \\citealt{Condon_etal.1998}) and the Faint Images of the Radio Sky at Twenty centimeters~(FIRST; \\citealt{Becker_etal.1995}) at 1.4~GHz.  Because of the resolution effect due to the smaller beam size of FIRST, we could find an intrinsic variability only if the flux density in FIRST was larger than that in NVSS.  Data for all sources except for RX~J0806.6+7248 were available from both NVSS and FIRST.  Only FBQS~J1629+4007 showed a FIRST flux density larger than that in NVSS.  Using the same technique as \\citet{Ghosh&Punsly2007}, we detected a flux variation in FBQS~J1629+4007 with a significance of $5.6\\sigma$; the brightness temperatures were derived to be $10^{13.0}$~K~(Table~\\ref{table5}).  Based on the near equipartition of energies between radiating particles and a magnetic field \\citep{Readhead1994} or the catastrophic electron energy loss due to an inverse Compton process \\citep{Kellermann_Pauliny-Toth1969}, the brightness temperature of nonthermal radio emission is expected to be within a physically realistic upper limit of $10^{11.2}$~K or $10^{12}$~K in the rest frame; the Doppler factor is expected to be $\\delta>4.0$ or 2.1 and the viewing angles should be 14$\\arcdeg$ or 28$\\arcdeg$, respectively.  Thus, it is evident that the emitting source is a highly relativistic jet viewed pole-on.  This source is presumably Doppler boosted and may be intrinsically in the radio-quiet population~(Column~13 of Table~\\ref{table5}).  The variability Doppler factors for several objects have also been previously reported in a very radio-loud NLS1 population \\citep{Yuan_etal.2008}.  Our limited dynamic-range VLBA image of FBQS~J1629+4007 showed only a structure of two discrete components (the Appendix), which probably indicated a one-sided jet.  It is important to conduct VLBI imaging more thoroughly and at more observing frequencies in the future.     \\subsection{Nature of Radio-loud NLS1 Radio Sources}\\label{section:naturesofRLNLS1s} NLS1s are expected to be radio-quiet AGNs on the basis of the possible dependence of radio loudness on both the black hole mass and the accretion rate~\\citep{Boroson&Green1992}.  Even if all NLS1s were intrinsically radio quiet, the Doppler beaming effect could explain why some NLS1s appeared radio loud.  Our VLBA study for the five radio-loud NLS1s has shown both the cases that are likely and unlikely due to the Doppler beaming effect.  Although they may be intrinsically radio quiet, FBQS~J1629+4007 and FBQS~J1644+2619 (inverted-spectrum sources) have shown evidence of a significant Doppler beaming effect, which presumably contributes to their radio loudness.  We emphasize that the strong Doppler beaming requires not only a nearly pole-on viewed orientation but also a relativistic bulk jet speed.  The nuclei of these NLS1s are expected to have the ability to accelerate the jets to the relativistic speed.  In the projected distance, the asymmetry of FBQS~J1644+2619 continues up to 30~kpc; kpc-scale asymmetry has also been seen in many blazars in their VLA images \\citep{Murphy_etal.1993,Kharb_etal.2010}.  The jet length 30~kpc corresponds to $\\sim 4 \\times 10^{10}$~the Schwarzschild radius for FBQS~J1644+2619 (an estimated black hole mass of $8.4 \\times 10^6\\ M_\\mathrm{\\sun}$; \\citealt{Yuan_etal.2008}).  The largest radio galaxies have radio structures of $>1$~Mpc \\citep{Palma_etal.2000,Machalski_etal.2008}, which corresponds to $>1 \\times 10^{10}$~the Schwarzschild radius of a black hole mass of $10^9\\ M_\\mathrm{\\sun}$.  The significance of FBQS~J1644+2619 is that at least some NLS1 nuclei have the ability to be huge radio galaxies in black-hole-mass-normalized scales.  A further study of the kpc-scale jets of FBQS~J1644+2619 will be reported elsewhere.  RX~J0806.6+7248 and B3~1702+457, which are steep-spectrum sources with non-negligible diffuse structures, are unlikely to be strongly beamed; this indicates intrinsic radio-loud sources without Doppler boosting.  RX~J0806.6+7248 shows jet bending at $\\sim$20~pc from the core; this jet bending extends to $\\sim$100~pc~(Figure~\\ref{figure1}).  Such a structure resembles that of TeV gamma-ray blazar Mrk~501, which has inner jets strongly bending at $\\sim30$~mas from the core and continuing onto very diffuse outer jets.  \\citet{Giroletti_etal.2004} suggested that highly relativistic bulk speeds and a progressive change in the jet direction ($4$--$25\\degr$) with respect to the line of sight are possible and in agreement with the observed one-sided jet properties in the pc scale of Mrk~501.  A quite small viewing angle is required to make such an intrinsically small bending appear large.  The core spectrum inverted up to 8.4~GHz \\citep{Giroletti_etal.2004} can attribute to a strong beaming effect on the spectral peak frequency $\\nu_\\mathrm{obs} = \\delta \\nu_\\mathrm{rest}$ (Section~\\ref{section:Dopplerbeamingeffect}) for pole-on viewed innermost jets.  On the other hand, the inner jets of RX~J0806.6+7248 appear two-sided with respect to the core with a steep spectrum, and the extended jets originate from the both sides of the two-sided structure (the Appendix); the extended jets of Mrk~501 distinctly originate from the one-sized inner jets.  It is difficult to attribute these radio properties of RX~J0806.6+7248 to beamed jets.   Thus, the nuclei of some NLS1s can generate jets that are strong enough to make them intrinsically radio loud.  Consequently, either case (strongly or weakly beamed) suggests significant jet productivity in the radio-loud NLS1s in our sample.  NLS1s as a class show a number of extreme properties and radio-loud ones are very rare.  However, we build on these radio results to understand that the central engines of radio-loud NLS1s are essentially the same as that of other radio-loud AGNs in terms of the formation of nonthermal jets.   NLS1s show two types of X-ray continuum spectra.  Normal~(radio-quiet) NLS1s show a prominent soft excess plus a relatively steep ($\\Gamma \\sim 2$) power-law hard tail \\citep{Haba_etal.2008,Leighly1999b}.  Both the components are proposed to originate from the accretion disks; the soft excess is thermal emission from the extremely slim disk \\citep{Abramowicz_etal.1988,Mineshige_etal.2000,Wang&Netzer2003} and the power-law hard component originates from the scattered-up disk photons generated by the inverse Compton process.  On the other hand, \\citet{Yuan_etal.2008} recently indicated that while some radio-loud NLS1s show X-ray spectra similar to normal NLS1s, many radio-loud NLS1s resemble the high-energy-peaked flat-spectrum radio quasars~\\citep[HFSRQs;][]{Perlman_etal.1998} in terms of the broadband spectral energy distribution~(SED), which shows synchrotron spectra (due to beamed jets) peaked at UV/X-ray regimes, resulting in very steep X-ray spectra.  FBQS J1629+4007 and FBQS~J1644+2619, both of which are presumably Doppler beamed sources in our study, show HFSRQ-like SEDs \\citep{Padovani_etal.2002,Yuan_etal.2008}.  In particular, FBQS~J1629+4007 is the first confirmed HFSRQ for its modeled synchrotron peak at $2\\times10^{16}$~Hz and its very steep X-ray spectrum.  \\citet{Zhou_etal.2007} also suggested that the radio-loud NLS1 J0324+3410 was an NLS1-blazar composite with an HFSRQ-like SED.  Although a small number of sources in our study are discussed here, the radio-loud NLS1s with beamed jets may have jet-dominated HFSRQ-like SEDs.  On the other hand, RX~J1633.3+4718, which may be unlikely to be strongly Doppler beamed because of its steep spectrum, shows an X-ray spectrum with a soft excess plus a power-law hard tail similar to several normal NLS1s \\citep{Yuan_etal.2008}.  \\citet{Yuan_etal.2010} showed this ultra-steep soft X-ray emission can be well fitted by an optically thick accretion disk, whose inferred parameters (a black hole mass of $\\sim 3 \\times 10^6\\ M_\\mathrm{\\sun}$ and accretion rate) are in good agreement with the independent estimates based on an optical emission line spectrum.  \\citet{Yuan_etal.2008} discussed that the hard tail is flat~($\\Gamma=1.37 \\pm 0.49$) and may be an inverse Compton component originating from the jets, similar to the flat-spectrum radio quasars.  \\citet{Abdo_etal.2009a} also discussed that the simultaneous optical/X-ray/$\\gamma$-ray SED of the very radio-loud NLS1 SDSS~J094857.31+002225.4 showed a blue continuum attributed to the accretion disk and a hard X-ray spectrum attributed to the jet.  Thus, the contribution of the disk could not be negligible even in the radio-loud NLS1s.  Perhaps, even the HFSRQ-like SED may be due to only the dominance of the thermal emission from the extremely slim disk.  We would have to compare many SEDs and beamed-jet profiles by conducting more VLBI imaging observations for radio-loud NLS1s in order to understand the origin of the X-ray spectra.  Similar to NLS1s, broad absorption line~(BAL) quasars \\citep{Weymann_etal.1991} are also considered to be highly accreting black hole (with larger masses compared to NLS1s) systems \\citep{Boroson2002} and generally radio quiet \\citep{Becker_etal.2001}.  The rarity of the BAL quasars as radio galaxies \\citep{Gregg_etal.2006} is an important similarity to NLS1s.  A VLBI study of BAL quasars \\citep{Jiang_Wang2003,Montenegro-Montes_etal.2009,Doi_etal.2009,Kunert-Bajraszewska_etal.2010} is crucial for revealing their accreting and outflowing phenomena and determining their speed and orientation of the nonthermal jets by observing them.  An understanding of the natures of the highly accreting systems is required to explain both the subclasses of NLS1s and BAL quasars.    The VLBA observations at 1.7~GHz by direct imaging of pc-scale radio structure have provided convincing evidence of nonthermal jets in radio-loud NLS1s.  We resolved significant pc-scale structures of three NLS1s; the pc-scale radio emissions provided most of the total fluxes found in the VLA resolutions in all the targets.  The previously reported JVN images at 8.4~GHz enabled us to determine that our sample contained two inverted- and three steep-spectrum sources.  FBQS~J1644+2619, an inverted spectrum source, showed a remarkable one-sided pc-scale jet structure that indicated either a mildly or highly relativistic jet speed of $\\beta>0.74$ and an inclination of $<42\\degr$ on the basis of the framework of the Doppler beaming effect.  FBQS~J1629+4007, an inverted-spectrum source, showed compact radio structures in the pc-scale and a significant flux variability between the data from NVSS and FIRST.  The physically unrealistic brightness temperature inferred from the variability requires a pole-on viewed relativistic jet of a Doppler factor of $\\delta>4.0$, implying a jet speed of $\\beta>0.88$ and an inclination of $<14\\degr$.  The Doppler beaming effect is possibly responsible for the radio loudness in FBQS~J1644+2619 and FBQS~J1629+4007.  RX~J0806.6+7248 and B3~1702+457, which are steep-spectrum sources with non-negligible diffuse structures in the pc scale, were unlikely to be strongly beamed; this indicates that they are intrinsically radio loud even without the beaming effect.  We found both types of radio-loud NLS1s with relativistically beamed jets and intrinsically strong jets, as well as seen in other radio-loud AGN classes."
    },
    "1107/1107.1804_arXiv.txt": {
        "abstract": "We investigate the effects of dust on Ly$\\alpha$ photons emergent from an optically thick medium by solving the integro-differential equation of the radiative transfer of resonant photons. To solve the differential equations numerically we use the Weighted Essentially Non-oscillatory method (WENO). Although the effects of dust on radiative transfer is well known, the resonant scattering of Ly$\\alpha$ photons makes the problem non-trivial. For instance, if the medium has the optical depth of dust absorption and scattering to be $\\tau_a \\gg 1$, $\\tau \\gg 1$, and $\\tau \\gg \\tau_a $, the effective absorption optical depth in a random walk scenario would be equal to $\\sqrt{\\tau_a(\\tau_a+\\tau)}$. We show, however, that for a resonant scattering at frequency $\\nu_0$, the effective absorption optical depth would be even larger than $\\tau(\\nu_0)$. If the cross section of dust scattering and absorption is frequency-independent, the double-peaked structure of the frequency profile given by the resonant scattering is basically dust-independent. That is, dust causes neither narrowing nor widening of the width of the double peaked profile. One more result is that the time scales of the Ly$\\alpha$ photon transfer in the optically thick halo are also basically independent of the dust scattering, even when the scattering is anisotropic.  This is because those time scales are mainly determined by the transfer in the frequency space, while dust scattering, either isotropic or anisotropic, does not affect the behavior of the transfer in the frequency space when the cross section of scattering is wavelength-independent. This result does not support the speculation that dust will lead to the smoothing of the brightness distribution of Ly$\\alpha$ photon source with optical thick halo. ",
        "introduction": "Ly$\\alpha$ photons have been widely applied to study the physics of luminous objects at various epochs of the universe, such as Ly$\\alpha$ emitters, Ly$\\alpha$ blob, damped Ly$\\alpha$ system, Ly$\\alpha$ forest, fluorescent Ly$\\alpha$ emission, star-forming galaxies, quasars at high redshifts as well as optical afterglow of gamma ray bursts (Haiman et al. 2000; Fardal et al. 2001; Dijkstra \\& Loeb. 2009; Latif et al. 2011). The resonant scattering of Ly$\\alpha$ photons with neutral hydrogen atoms has a profound effect on the time, space and frequency dependencies of Ly$\\alpha$ photons transfer in an optically thick medium. Ly$\\alpha$ photons emergent from an optically thick medium would carry rich information of photon sources and halo surrounding the source of the Ly$\\alpha$ photon. The profiles of the emission and absorption of the Ly$\\alpha$ radiation are powerful tools to constrain the mass density, velocity, temperature and the fraction of neutral hydrogen of the optically thick medium. Radiation transfer of Ly$\\alpha$ photons in an optically thick medium is fundamentally important.  The radiative transfer of Ly$\\alpha$ photons in a medium consisting of neutral hydrogen atoms has been extensively studied either analytically or numerically. Yet, there have been relatively few results which are directly based on the solutions of the integro-differential equation of the resonant radiative transfer. Besides the Field solution (Field 1959, Rybicki \\& Dell'Antonio 1994), analytical solutions with and without dust mostly are based on the Fokker-Planck (P-F) approximation (Harrington 1973, Neufeld 1990, Dijkstra et al. 2006). The P-F equation might miss the detailed balance relationship of resonant scattering (Rybicki 2006), and therefore, the analytical solutions cannot describe the formation and evolution of the Wouthuysen-Field (W-F) local thermalization of the Ly$\\alpha$ photon frequency distribution (Wouthuysen 1952, Field 1958), which is important for the emission and absorption of the hydrogen 21 cm line (e.g. Fang 2009). The features of the Ly$\\alpha$ photon transfer related to the W-F local thermalization are also missed. An early effort (Adams et al. 1971) trying to directly solve the integro-differential equation of the resonant radiative transfer with numerical method. It still is, however, of a time-independent approximation.  Recently, a state-of-the-art numerical method has been introduced to solve the integro-differential equation of the radiative transfer with resonant scattering (Qiu et al. 2006, 2007, 2008, Roy et al. 2009a). The solver is based on the weighted essentially non-oscillatory (WENO) scheme (Jiang \\& Shu 1996). With the WENO solver, many physical features of the transfer of Ly$\\alpha$ photons in an optically thick medium (Roy et al. 2009b, 2009c, 2010), which are missed in the Fokker-Planck equation approximations, have been revealed. For instance, the WENO solution shows that the time scale of the formation of the W-F local thermal equilibrium actually is only about a few hundred times of the resonant scattering. It also shows that the double peaked frequency profile of the Ly$\\alpha$ photon emergent from an optically thick medium does not follow the time-independent solutions of the P-F equation. These results directly indicate the needs of re-visiting problems which have been studied only via the F-P time-independent approximation.  We will investigate, in this paper, the effects of the dust on the Ly$\\alpha$ photons transfer in an optically thick medium. Dust can be produced at epochs of low and moderate redshifts, and even at redshift as high as 6 (Stratta et al. 2007). Absorption and scattering of dust have been used to explain the observations on Ly$\\alpha$ emission and absorption (Hummer \\& Kunasz 1980), such as the escaping fraction of Ly$\\alpha$ photons (Hayes et al. 2010, 2011, Blanc et al. 2010); the redshift-dependence of the ratio between Ly$\\alpha$ emitters and Lyman Break galaxies (Verhamme et al. 2008); and the ``evolution'' of the double-peaked profile (Laursen et al. 2009).  Nevertheless, it is still unclear whether the time scale of photon escaping from optically thick halo will be increasing (or decreasing) when the halo is dusty. It is also unclear whether the effects of dust absorption can be estimated by the random walk picture (Hansen \\& Oh 2006). As for the dust effect on the double-peaked profile, the current results given by different studies seem to be contradictory: some claims that the dust absorption leads to the narrowing of the double-peaked profile (Lauresen et al 2009), while others result that the width between the two peaks apparently should be increasing due to the dust absorption (Verhamme et al. 2006). We will focus on these basic problems, and examine them with the solution of the integro-differential equation of radiative transfer.  This paper is organized in the following way: section 2 presents the theory of the Ly$\\alpha$ photon transfer in an optically thick medium with dust. The equations of the intensity and flux of resonant photons in a dusty medium are given. We will study three models of the interaction between dust and photons:  (1)  dust causes only scattering with photons; (2) dust causes both scattering and absorption; and (3) dust causes only absorption of photons. Section 3 gives the solutions of Ly$\\alpha$ photons escaping from an optically thick spherical halos with dust. The dusty effect on the double-peaked profile will be studied in Section 4. The discussion and conclusion are given in Section 5. Some mathematical derivations of the equations and numerical implementation details are given in the Appendix. ",
        "conclusions": "The study of dust effects on radiative transfer has had a long history related to extinction. However, dust effects on radiative transfer of resonant photons actually have not been carefully investigated. Existing works are mostly based on the solutions of the Fokker-Planck approximation, or Monte Carlo simulation. These results are important. We revisited these problems with the WENO solver of the integro-differential equation of the resonant radiative transfer, and have found some features which have not been addressed in previous works.  These features are summarized as follows.  First, the random walk picture in the physical space will no longer be available for estimating the effective optical depth of dust absorption. For a medium with the optical depth of absorption and resonant scattering to be $\\tau_a\\gg 1$, $\\tau(\\nu_0)\\gg 1$ and $\\tau_s(\\nu_0)\\gg \\tau_a$, the effective absorption optical depth is found to be almost independent of $\\tau_a$, and to be equal to about a few times of $\\tau_s(\\nu_0)$.  Second, dust absorption will, of course, yield the decrease of the flux of Ly$\\alpha$ photons emergent from optical thick medium. However, if the absorption cross-section of dust is frequency independent, the double-peaked structure of the frequency profile is basically dust-independent. The double-peaked structure does not get narrowed or widened by the absorption and scattering of dust.  Third, the time scales of Ly$\\alpha$ photon transfer basically are independent of dust scattering and absorption. It is because those time scales are mainly determined by the kinetics in the frequency space, while dust does not affect the behavior of the transfer in the frequency space if the cross section of the dust is wavelength-independent. The local thermal equilibrium makes the anisotropic scattering to be ineffective on the angular distribution of photons. Dust absorption and scattering do not lead to the increase or decrease of the time of storing Ly$\\alpha$ photons in the halos.  The differences between the time-independent solutions of the Fokker-Planck approximation, or Monte Carlo simulation and the WENO solution of eq.(3) is mainly related to the W-F effect. Therefore, all above-mentioned features can already be clearly seen with halos of $\\tau_0 \\sim 10^2$, in which the W-F local thermal equilibrium has been well established.  In this context, most calculation in this paper is on holes with $\\tau_0<10^5$. This range of $\\tau_0$ certainly is unable to describe halos with column number density of HI larger than 10$^{17}$ cm$^{-2}$ (e.g. Roy et al. 2010). Nevertheless, the result of $\\tau_0<10^5$ would already be useful  for studying the 21 cm region around high-redshift sources, of which the optical depth typically is (Liu et al 2007; Roy, et al. 2009c). \\begin{equation} \\tau_0=3.9\\times 10^5 f_{\\rm HI}\\left (\\frac{T}{10^4 {\\rm K}}\\right )^{-1/2} \\left(\\frac{1+z}{10}\\right)^{3}\\left ( \\frac{\\Omega_bh^2}{0.022}\\right ) \\left ( \\frac{R_{\\rm ph}}{10 {\\rm kpc}}\\right ), \\end{equation} where $f_{\\rm HI}$ is the fraction of HI. All other parameters in eq. (15) is taken from the concordance $\\Lambda$CDM mode. For these objects the relation between dimensionless $\\eta$ and physical time $t$ is given by \\begin{equation} t=5.4 \\times 10^{-2}f^{-1}_{\\rm HI} \\left (\\frac{T}{10^4 {\\rm K}}\\right )^{1/2} \\left(\\frac{1+z}{10}\\right)^{-3}\\left ( \\frac{\\Omega_bh^2}{0.022}\\right )^{-1} \\eta, \\ \\ {\\rm yr}. \\end{equation} The 21 cm emission rely on the W-F effect. On the other hand, the time-scale of the evolution of the 21 region is short. The effect of dust on the time-scales of Ly$\\alpha$ evolution should be considered.  We have not considered the Ly$\\alpha$ photons produced by the recombination in the ionized halo. If the halo is optical thick, photons from the recombination will also be thermalized. The information of where the photon comes from will be forgotten during the thermalization. Therefore, photons from recombination should not show any difference from those emitted from central sources. Only the photons formed at the place very close to the boundary of the halo will not be thermalized, and may yield different behavior."
    },
    "1107/1107.3801_arXiv.txt": {
        "abstract": "Studying the average properties of active galactic nuclei (AGN) host stellar populations is an important step in understanding the role of AGN in galaxy evolution and the processes which trigger and fuel AGN activity.  Here we calculate model spectral energy distributions (SEDs) that include emission from the AGN, the host galaxy stellar population, and dust enshrouded star formation.  Using the framework of cosmic X-ray background population synthesis modeling, the model AGN hosts are constrained using optical (B band) and near infrared (J band, 3.6 $\\mu$m, 5.7 $\\mu$m, 8.0 $\\mu$m, and 24 $\\mu$m) luminosity functions and number counts.  It is found that at $z$ $<$ 1, type 1 and type 2 AGN hosts have similar stellar populations, in agreement with the orientation based unified model and indicative of secular evolution.  At $z$ $>$ 1, type 2 AGN hosts are intrinsically different from type 1 AGN hosts, suggesting that the simple orientation based unified model does not hold at $z$ $>$ 1. Also, it is found that if Compton thick (CT) AGN evolve like less obscured type 2 AGN, then, on average, CT AGN hosts are similar to type 2 AGN hosts; however, if CT obscuration is connected to an evolutionary stage of black hole growth, then CT AGN hosts will also be in specific evolutionary stages. Multi-wavelength selection criteria of CT AGN are discussed. ",
        "introduction": "\\label{sect:intro}  It is known that all massive galaxies have a central supermassive black hole (SMBH) \\citep[e.g.,][]{KR95}.  Yet it is unknown why only a small fraction of these SMBHs are actively accreting as active galactic nuclei (AGN). Several theories attempt to explain the fueling mechanism of AGN: major mergers \\citep[e.g.,][]{S88, H06}, minor mergers and gravitational instabilities within the host galaxy \\citep[e.g.,][]{Cr03,KK04, P07,S08}, nuclear starbursts \\citep[e.g.,][]{D07,B08}, supernova explosions within nuclear starbursts \\citep[e.g.,][]{C09, KJ10}, and collisions of warm halo clouds with the nuclear region \\citep{M10}.  Given the large range of observed AGN properties, it is likely that different fueling mechanisms come into play for different AGN populations.  Current observations suggest that powerful quasars are triggered by major mergers but moderate luminosity AGN are more likely to be triggered by stochastic fueling incidents \\citep{B06a,H08,HH09,L10}.  However, it is not yet clear which fueling mechanisms are dominant in which portions of the AGN population.  The unified model of AGN \\citep{A93} explains the different observed levels of AGN obscuration as a simple geometric effect.  This model assumes that the central engines of all AGN are identical and the level of obscuration is dependent on the line of sight between the observer and the central engine.  Thus type 2 AGN are obscured by column densities $N_H$ $\\gtrsim$ 10$^{22}$ cm$^{-2}$, because the observer is looking through the dusty torus, in contrast to type 1 AGN which are unobscured ($N_H$ $<$ 10$^{22}$ cm$^{-2}$) because the observer is looking down the throat of the dusty torus.  Since the central engines are identical, it is expected that, on average, the host galaxies of various spectral types of AGN will be similar.  However, different black hole fueling mechanisms are likely to lead to different relationships between the various spectral types of AGN.  In the major merger AGN fueling paradigm, the AGN is triggered while deeply embedded in gas and dust and thus when the AGN first turns on, it is highly obscured.  Eventually the radiation pressure pushes away the remaining gas and dust revealing an unobscured quasar \\citep[e.g.,][]{P04,H06,Ri09}.  In the major merger paradigm, type 1 and type 2 AGN have host galaxies which are in different evolutionary stages. Conversely, in the nuclear starburst fueling paradigm, type 1 and type 2 AGN hosts can be similar since on average the unified model can hold \\citep{B08}. This shows that tests of the unified model can be used to explore which processes are viable options for AGN fueling and in which sections of the AGN population different mechanisms are relevant.  Classifying the host galaxies of Compton thick (CT) AGN, AGN with $N_H$ $\\gtrsim$ 10$^{24}$ cm$^{-2}$, is of special interest. It is expected that a large fraction of AGN with CT levels of obscuration are in a young evolutionary stage characterized by rapid black hole and host bulge growth \\citep{S88,F99,H06,F09,DB10,T10}.  Thus, understanding the host properties of CT AGN may offer special insight into the AGN triggering process and the formation of massive galaxies.  Large samples of CT AGN are difficult to identify because, by definition, CT AGN suffer from extreme levels of obscuration \\citep[e.g.,][]{Ghis94}.  Therefore, previous studies of AGN host galaxies have not been able to study the stellar populations of CT AGN.  Galaxy optical colors are often described in terms of the red sequence and the blue cloud, where objects located on the red sequence are characterized by massive older stellar populations and objects on the blue cloud are less massive and have young stellar populations and current star formation.  Between the red sequence and the blue cloud is a lightly populated region referred to as the green valley.  Some studies find that AGN feedback is connected to the shut down of host star formation \\citep[e.g.,][] {Mc10, B11} and thus certain types of AGN hosts preferentially reside in the green valley \\citep[e.g.,][]{H09,GS10,W10,SR11}.  Thus AGN activity is potentially a stage in galaxy evolution where feedback from the AGN could shut down star formation, causing the host galaxy to age from the blue cloud, across the green valley, and onto the red sequence \\citep{F07,S07}.  However, other studies find that the colors of an AGN and its host are more closely related to the amount of available obscuring material rather than the evolutionary stage of the host stellar population \\citep{B09,Geo09,C10,R11}.  Indeed, \\citet{C10} show that AGN are intrinsically bimodal in color and that many AGN which appear to be on the red sequence are actually located in the blue cloud once dust extinction is taken into account.  Thus further investigation is necessary to establish the nature of AGN host galaxy stellar populations and the role of AGN feedback in regulating star formation.  This point is especially salient as several possible AGN fueling mechanisms include either causal or concurrent star formation. Some studies even show that AGN obscuration may in part be due to nuclear starburst disks \\citep[e.g.,][]{T05, D07,B08}.  In order to understand how host galaxy processes affect the central SMBH and the role of the central SMBH in host galaxy evolution, the nature of stellar populations of AGN hosts must be well understood.  As AGN hosts exhibit size-able object-to-object variability, it is difficult to elucidate trends in the AGN population by fitting individual objects.  Thus, in order to study larger trends in AGN--host galaxy interactions, the average stellar properties of a large ensemble of AGN hosts is investigated here.  To this end the stellar synthesis models of \\citet{bc03} are used to explore the stellar populations of AGN hosts. As in previous studies \\citep{B06,DB11}, the AGN SEDs are calculated using the photoionization code Cloudy \\citep{F98}.  The cosmic X-ray background (CXRB) synthesis modeling framework is used to characterize the AGN population.  The model SEDs are then used to move the CXRB framework into other wavelength regions.  This allows the space density and evolution of AGN host galaxies to be determined by the most comprehensive census of AGN activity. By comparing these models against various observations in the optical through mid-infrared (mid-IR) spectral ranges, constraints are placed on the average stellar populations of type 1, Compton thin type 2 (here referred to simply as type 2), and CT AGN at various redshifts.  In Section \\ref{sect:model} the AGN spectral model and stellar population model are described. Section \\ref{sect:results} presents the results of the model.  In Sections \\ref{sect:disc} and \\ref{sect:sum} the results are discussed and summarized.  AB magnitudes are used, unless otherwise stated, and $h$ = 0.7, $\\Omega_{\\Lambda}$ = $1-\\Omega_{M}$ = 0.7 is assumed as necessary. ",
        "conclusions": "\\label{sect:disc}  \\subsection{Hosts of CT AGN} \\label{sub:CT}  For both the evolving and non-evolving models, the CT AGN host stellar populations suffer from at least as much dust extinction as the type 2 AGN host stellar populations. Observations suggest that at least some of the observed extinction of heavily obscured AGN may be due to extended dust structures or molecular clouds within the host galaxy \\citep{B07, P08, MS10}, so it is expected that the stellar populations of CT AGN will be enshrouded in dust.  It is within the limits imposed by the model constraints for the high Eddington ratio CT AGN hosts to have an average $E(B-V)$ which is in agreement with ULIRGs hosting AGN \\citep{RZ10}.  The AGN evolution scheme, in which mergers trigger large nuclear starbursts and AGN activity, claims that the AGN activity will initially be very highly obscured while the black hole grows very rapidly.  As the black hole grows, the radiation pressure on surrounding dusty gas will increase until the AGN feedback blows out the obscuring material and halts the star formation in the host nuclear region \\citep{S88, P04, Ri09, H06}.  In this scheme CT AGN should have young stellar populations and possibly high levels of on going star formation.  The non-evolving model places a lower limit on the stellar population age of 1 Gyr.  However, when the enhanced star formation sources are included, the non-evolving model CT AGN hosts have similar stellar ages to the type 2 AGN hosts, in contrast to the expectations of the AGN evolution scheme.  For the high Eddington ratio CT AGN in the evolving model, the lower limit on the stellar population age is 0.3 Gyr, the average stellar age of ULIRGs hosting an AGN \\citep{RZ10}.  The low Eddington ratio CT AGN hosts of the evolving model can be of similar age as the type 1 and type 2 AGN, but the near and mid-IR predicted number counts are in better agreement with the observations if these AGN hosts have slightly older stellar populations than the average type 1 and type 2 AGN hosts. It has been demonstrated that galaxies hosting AGN with lower [OIII] luminosities have larger values for $D_n$(4000 \\AA) compared to galaxies hosting AGN with high [OIII] luminosities \\citep{K03,K04,K07}. Thus it is expected that AGN with lower Eddington ratios are in hosts with older stellar populations, in agreement with the findings of this study.  The stellar populations of CT AGN hosts for the evolving model and the non-evolving model are in agreement with the expectations from the AGN evolution scheme.  The non-evolving model finds that the average CT AGN host has recently ($\\gtrsim$1 Gyr ago) undergone a large burst of star formation, but that current star formation rates are more modest.  The evolving model finds that high Eddington ratio CT AGN hosts have recently ($\\gtrsim$0.3 Gyr ago) undergone a large burst of star formation and that current star formation rates may also be elevated.  If the enhanced star formation sources are included, the non-evolving model AGN hosts have stellar populations of similar age as the type 2 AGN, and thus are in better agreement with the orientation based unified model than with the AGN evolution scheme.  The evolving model CT AGN hosts are consistent with the AGN evolution scheme regardless of the inclusion of the enhanced star formation sources.  In summary, the non-evolving model CT AGN hosts are a simple extension of the type 2 AGN host population while the evolving model CT AGN are consistent with the paradigm where major mergers cause both intense starbursts and AGN activity.  In the evolving model, low Eddington ratio CT AGN hosts will appear as galaxies with old and dusty stellar populations while the high Eddington ratio CT AGN hosts will appear as IR bright starburst galaxies.   \\subsection{Enhanced Star Formation in AGN Hosts} \\label{sub:ensf}  It is found that $\\sim$5-15$\\%$ of AGN hosts can be enhanced star formation sources with an average SFR $\\approx$ 20 M$_{\\odot}$ yr$^{-1}$, a factor of 10 higher than the majority of AGN hosts.  This is expected to be a lower limit of the fraction of enhanced star formation sources as some sources may have such highly embedded star formation that the majority of the reprocessed emission due to star formation is at wavelengths longer than 24 $\\mu$m.  Observations at longer wavelengths, such as in the far-IR with the {\\em Herschel} or the millimeter/sub-mm regime with the Atacama Large Millimeter/submillimeter Array (ALMA), are necessary for uncovering the evolution of star formation rates in AGN hosts \\citep{DB11}.  Even deep radio observations are a useful tool in determining the highly embedded star formation rates of AGN hosts \\citep{Bal09}. Indeed, by stacking sub-mm observations of X-ray selected AGN, \\citet{L10} found that the average AGN host SFR $\\approx$ 30 M$_{\\odot}$ yr$^{-1}$. Furthermore, \\citet{L10} found  that the AGN host SFR evolves strongly with redshift but with evidence of luminosity dependent evolution only for the highest luminosity AGN.  In the scenario where AGN host SFRs are not luminosity dependent, it is found that more than half of AGN hosts can be enhanced star formation sources.   The fact that the mid-IR AGN and host number counts are over-predicted by an average SFR $>$ 2 M$_{\\odot}$ yr$^{-1}$, despite observational evidence that there is a population of AGN with an average SFR an order of magnitude higher than this upper limit, suggests that there are two populations of AGN.  The majority of AGN hosts, at least at $z$ $<$ 1, have SFRs similar to local spiral galaxies and $\\gtrsim$5-15$\\%$ of AGN hosts have markedly higher SFRs.  The existence of these two populations of AGN hosts can be interpreted in two complimentary ways.  The first interpretation is that the AGN hosted by galaxies with enhanced SFRs are being fueled by different mechanisms than the lower star formation objects. At $z$ $<$ 1, the lower star formation objects are likely dying quasars which were triggered by major mergers while the enhanced star formation objects are likely Seyferts which are both fueled and obscured by circumnuclear starburst disks, as investigated by \\citet{B08}.  The other interpretation is that all AGN are fueled by processes related to nuclear starbursts, whether those starbursts are triggered by mergers or through secular processes, and on average, the nuclear starburst phase overlaps with the active AGN phase for $\\gtrsim$5-15$\\%$ of the AGN lifetime.  Further exploration of these two populations of AGN hosts is necessary to determine the processes which trigger and fuel AGN activity.    \\subsection{Methods for Finding CT AGN} \\label{sub:findCT}   It is well documented that the integrated emission of AGN observed in deep X-ray surveys is insufficient to account for the intensity of the CXRB at $\\sim$30 keV and that the shape of the CXRB necessitates that the missing population of AGN be highly obscured \\citep[e.g.,][]{B06a, DB09, T09b}. Given the uncertainties of the normalization of the CXRB and the AGN hard X-ray luminosity functions, different models predict vastly different numbers of missing CT AGN \\citep[e.g.,][]{G07,DB09,T09b}. As these highly obscured AGN are missed in deep X-ray surveys \\citep{H08}, it is important to consider other methods to identify the elusive CT AGN population.  One possibility is that the majority of X-ray bright optically inactive galaxies (XBONGs) host CT AGN \\citep{Fi08,Fi09,T09,R10}.  XBONGs are X-ray sources found in deep surveys which have no optical counterparts with $R$ $\\lesssim$ 25.5 and make up a substantial portion of deep survey X-ray sources \\citep{A10}.  It is thought that these sources are either heavily obscured AGN and/or high redshift quasars \\citep{A01,M05,R10}.   Considering the sources above the dotted horizontal line in Figures \\ref{fig:fxtororig} and \\ref{fig:fxtorcomp} suggests that a small fraction of optically faint X-ray sources may be CT AGN, but the majority of this population is $z$ $\\gtrsim$ 1 type 2 AGN.  Thus it is unlikely that XBONGs host the majority of the missing CT AGN population.  Another method used to identify CT AGN candidates is to search for infrared bright sources which have an infrared excess, usually defined by $f_{24}/f_R$ $\\gtrsim$ 1000, where $f_{24}$ is the 24 $\\mu$m flux \\citep[e.g.][]{P06, A08, D08, Fi08, Fi09, T09}.  Some concern has been raised that this method will also pick out lower redshift type 2 AGN masquerading as high redshift CT AGN \\citep{G10}. We therefore investigate the population of AGN selected by the infrared excess criteria.  \\citet{Fi08} suggest that using the criteria $f_{24}/f_R$ $\\gtrsim$ 1000 and $R$ - $K$ $>$ 4.5 selects a distinct class of sources, the majority of which are CT AGN.  However, according to both the evolving and non-evolving model CT AGN and host SEDs developed here, the \\citet{Fi08} criteria selects low to moderately X-ray bright AGN at moderate to high redshift.  The vast majority of the selected AGN are obscured, but a significant fraction are still Compton thin.  Similarly, \\citet{D08} find that the \\citet{Fi08} criteria is likely to select dusty star forming templates and low X-ray flux AGN.  However, \\citet{D08} find no evidence that the selected AGN are CT.  A more effective method of identifying CT AGN candidates is based on the selection criteria of \\citet{P08}: $f_{24}/f_R$ $\\gtrsim$ 1000 and $f_{24}$ $>$ 1.0 mJy.  Furthermore, it appears that samples with a lower $f_{24}$ limit, such as  $f_{24}$ $>$ 700 $\\mu$Jy or even $f_{24}$ $>$ 550 $\\mu$Jy, also contain a large fraction of highly obscured AGN \\citep[][and references therein]{D10}.  Fainter samples of infrared bright galaxies are found to be predominately powered by star formation rather than by AGN \\citep[e.g.,][]{Pope08}.  Figure \\ref{fig:f24tor}a shows that representatives of all spectral types of AGN can be found with $f_{24}/f_R$ $\\gtrsim$ 1000 and $f_{24}$ $>$ 550 $\\mu$Jy.  These sources will be located at all redshifts. However, the majority of type 1 and type 2 AGN are bright in the soft X-ray. For both the evolving and non-evolving models, the vast majority of infrared excess AGN with $f_{24}$ $>$ 550 $\\mu$Jy and f$_{0.5-2}$ $<$ 10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ are CT AGN, as shown in Figure \\ref{fig:f24tor}b.  Figure \\ref{fig:f24torbyz}b, which shows the redshift distribution of the CT AGN with f$_{0.5-2}$ $<$ 10$^{-15}$ erg cm$^{-2}$ s$^{-1}$, illustrates that these X-ray faint CT AGN are located at all redshifts. For the evolving model the vast majority of these CT AGN are high Eddington ratio sources. This criteria will select a small number of type 2 AGN with enhanced star formation, but the majority of sources selected in this manner are indeed CT.  If an infrared excess AGN has $f_{0.5-2}$ $\\lesssim$ 10$^{-16}$ erg cm$^{-2}$ s$^{-1}$ and $f_{24}$ $>$ 550 $\\mu$Jy, than the AGN is CT and located at $z$ $\\lesssim$ 2.  This is true for both the evolving and non-evolving models. An interesting consequence of this is that a size-able population of $z$ $>$ 2 CT AGN should have 10$^{-15}$ $>$ $f_{0.5-2}$ $>$ 10$^{-16}$ erg s$^{-1}$ cm$^{-2}$.  According to the evolving model, nearly all of the $f_{0.5-2}$ $\\lesssim$ 10$^{-16}$ erg cm$^{-2}$ s$^{-1}$ sources are high Eddington ratio sources and according to the non-evolving model these low X-ray flux sources are nearly all enhanced star formation sources. It is therefore expected that X-ray stacking of IR bright sources with $f_{24}/f_R$ $\\gtrsim$ 1000 and low soft X-ray flux will yield a large fraction of the $z$ $\\lesssim$ 2 high luminosity CT AGN population. This shows that combining observations in multiple spectral regions, such as mid-IR, optical, and X-ray observations, is the most efficient way of identifying CT AGN candidates.  Identifying and characterizing the elusive CT AGN population is a necessary part of understanding the history of accretion and galaxy evolution.  However, due to the observational challenges of studying highly obscured sources, multi-wavelength investigations are necessary to identify and understand the nature of CT AGN.     \\subsection{Implications for the Unified Model and AGN Fueling Mechanisms}% \\label{sub:ave}  Several recent studies suggest that there are two distinct processes which lead to AGN activity, secular evolution and merger events \\citep[e.g.,][]{B06a,H08,HH09,L10}. In the latter paradigm, quasar activity is activated by galaxy mergers which cause gas and dust to be funneled into the nuclear region \\citep[e.g.,][]{S88,F99,P04,H06}, whereas moderate luminosity AGN are fueled by gravitational instabilities internal to the host galaxy or through minor interactions \\citep[e.g.,][]{Cr03,KK04,P07,S08}.  It is probable that if different forms of AGN activity are caused by contrasting fueling mechanisms, the relationship between various AGN spectral types may be different for the high and moderate luminosity populations.  This study finds that the average type 1 and type 2 AGN hosts have similar stellar populations and similar levels of dust attenuation at $z$ $<$ 1.  At $z$ $>$ 1 type 1 and type 2 AGN hosts appear to be fundamentally different.  Even though the high redshift type 1 and type 2 hosts have similar stellar populations, the type 1 AGN and host B band luminosity function requires that type 1 AGN hosts have significantly lower levels of dust extinction than what is required for type 2 AGN hosts to be in agreement with the observed near and mid-IR number counts.  At $z$ $>$ 1, type 1 AGN hosts are intrinsically less dusty than type 2 AGN hosts. This suggests that the orientation based unified model works well for describing the local Seyfert population, but may not be appropriate for the high redshift, high luminosity quasar population.  This conclusion is consistent with the findings of several other recent studies \\citep[e.g.,][]{B06,HH09,DB10,L10}.  For galaxies which are evolving secularly, the unified model appears to be an apt description.  However, the violent growth experienced by black holes and their host bulges during major merger events does not appear to fit into the orientation based unified scheme.  It is likely that different AGN fueling mechanisms will result in different relationships between AGN spectral types.  In the merger scenario, young AGN are highly obscured and old quasars are unobscured \\citep[e.g.,][]{S88, F99,P04,H06}.  Different spectral types of AGN fueled by secular processes related to nuclear starbursts are expected to be in agreement with the orientation based unified model due to the disky nature of the nuclear starburst which is likely to both fuel and obscure the AGN \\citep[e.g.,][]{B08}.  In the evolving model, it is found that mergers can play a strong role in fueling high $L_X$, $z$ $>$ 1 quasars.  However, in the non-evolving model, when the enhanced star formation sources are considered, it appears that the dominant fueling mechanism is not mergers.  Instead, the difference between spectral types may only be that some galaxies have less dust and gas than others at $z$ $>$ 1.  In the non-evolving $f_{CT}$ model, it appears that secular processes are the dominate AGN fueling mechanism at all redshift. For both the evolving and non-evolving model, at $z$ $<$ 1 the processes which lead to AGN activity are most likely secular.  However, the findings of this study suggest that the dominant AGN fueling process changes at $z$ $\\sim$ 1, since at $z$ $<$ 1 the average type 1 and type 2 AGN hosts are more similar than at higher redshift.  In order to understand the mechanisms which trigger and fuel AGN, it is important for future studies to pay careful attention to which fueling mechanisms are dominant in different subsets of the AGN population.  Thus, it has been shown that the relationship between different AGN spectral types is a helpful tool for understanding the dominant processes responsible for triggering and fueling AGN activity.  It is found that at $z$ $<$ 1 the orientation based unified model holds, suggesting that the dominant AGN fueling mechanisms at $z$ $<$ 1 are secular processes.  At $z$ $>$ 1, the orientation based unified model does not seem to hold.  The evolving CT AGN model suggests that at $z$ $>$ 1 mergers are an important AGN fueling mechanism; however, the non-evolving CT AGN model suggests that even at $z$ $>$ 1 the dominant AGN fueling mechanisms are secular processes and that mergers play only a minor role in fueling AGN at all redshifts.    \\subsection{$L_X$ and redshift Evolution of $f_2$} \\label{sub:evolf2}  Understanding the evolution of $f_2$ is an important step in understanding the nature of AGN obscuration and the interplay between an AGN and it's host galaxy.  If $f_2$ evolves with AGN $L_X$, this suggests that the obscuring material is close enough to the central engine that dust sublimation \\citep{Law91} and/or radiation pressure fed winds \\citep{KK94} affect portions of the obscuring material. If $f_2$ evolves with redshift this would indicate that the evolution of the obscuring material is somehow connected to the evolution of the host galaxy \\citep{B06}.  While most model constraints used in this study are not very sensitive to the evolution of $f_2$, the type 1 AGN and host B band luminosity function does prove to be a valuable test of the evolution of $f_2$.  This study finds that $f_2$ must evolve with AGN $L_X$ and that $f_2$ probably evolves with redshift.  If $f_2$ does not evolve with $L_X$, the predicted 1 $<$ $z$ $<$ 2 type 1 AGN and host B band luminosity functions are not in agreement with observations.  Using the host galaxy stellar population age and $E(B-V)$ as free parameters does not allow for an appropriate fit to the data.  The over-prediction of $m_B$ $\\sim$ 23 sources at 1 $<$ $z$ $<$ 2 is not due to observational bias as $m_B$ $\\sim$ 23 is significantly brighter than the depth of B band coverage accessible to surveys \\citep[e.g., B band 5$\\sigma$ level for COSMOS is $m_B$ = 26.7;][]{Scoville07}.  \\citet{LE10} argue that $f_2$ does not evolve with $L_X$ and that apparent evolution of $f_2$ with $L_X$ is due to X-ray observational biases. However, the $z$ $>$ 1 type 1 AGN and host B band luminosity function data used here is based on optical selection criteria \\citep{HS90,C04}.  The \\citet{DC96} type 1 AGN and host B band luminosity function does use a sample defined through X-ray selection, and the $z$ $<$ 0.4 predicted type 1 AGN and host B band luminosity function is in reasonable agreement with observation even if $f_2$ does not evolve with $L_X$.  This suggests that X-ray observational biases do not create an artificial evolution of $f_2$ with $L_X$. Therefore, at least for $z$ $>$ 1, $f_2$ must evolve with $L_X$.  While this study cannot rule out the possibility that $f_2$ is constant with redshift, our findings suggest that $f_2$ does evolve mildly with redshift.  While some studies suggest that $f_2$ must evolve with redshift \\citep[e.g.][] {LF05,B06a,B06,TU06}, the evidence is still tentative.  If AGN are fueled by different mechanisms, one might assume that the type 1/type 2 ratio will be different for the different mechanisms.  Therefore, if different mechanisms dominate during different epochs, the type 1/type 2 ratio should also include some redshift evolution.  It is possible that the slower secular evolution processes of Seyfert galaxies do not require evolution of $f_2$ with redshift or that the evolution is very mild, since this form of galaxy evolution is governed by stochastic processes.  However, in the quasar regime, where AGN fueling is initiated by major mergers, a stronger evolution of $f_2$ with redshift may be necessary.  Indeed, the type 1 AGN and host B band luminosity function which is least well fit by a non-evolving $f_2$ is the highest redshift bin considered, which is the redshift bin closest to the peak of quasar activity.  While this study does not conclusively show that $f_2$ evolves with redshift, the findings presented here suggest that $f_2$ does exhibit some evolution with redshift.  It is found that $f_2$ evolves with $L_X$ and $f_2$ is likely to evolve with redshift.  This suggests that the obscuring medium is close enough to the central engine to be affected by dust sublimation \\citep{Law91} and/or radiation pressure fueled winds \\citep{KK94}.  Also, the type 1 AGN and host B band luminosity function is an excellent X-ray independent tool to test the fraction of type 1 AGN predicted by X-ray observations."
    },
    "1107/1107.3176_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\medskip  The powerful radiation from high-mass protostars was the main obstacle for understanding the formation of massive stars (see e.g. \\citealt{1971Larson, 1974Kahn,1977Yorke, 1994Beech}; see also the recent review of \\citealt{2007Zinnecker}).  It was thought that the strong radiation pressure from the protostar would stop the accretion, hence preventing the formation of stars with masses over  $\\sim$10~M$_{\\sun}$. Although this theoretical  problem can be overcome (e.g., \\citealt{1987Wolfire,1989Nakano,  1996Jijina, 2002Yorke, 2000Norberg, 2003McKee, 2005Krumholz}), it  remains yet unclear how these stars are actually formed. There are several hypotheses for the formation process. Perhaps the most important ones are:  (1) large accretion rates (3 or 4 orders of magnitude greater than those observed in low-mass protostars) through circumstellar disks (\\citealt{1995Walmsley, 1996Jijina, 2003McKee, 2005Zhang,2007Banerjee}),  (2) competitive accretion between small protostars of the same cluster, that  could result in eventual mergers (\\citealt{1998Bonnell,2008Clarke,2011Moeckel,2011Baumgardt}), and   (3) accretion of ionized material, even after the formation of a compact HII region generated by the protostar (\\citealt{2002Keto, 2006Keto}).  \\begin{deluxetable*}{lcrlrrcc} \\tablewidth{0pc} \\tablecaption{Spectral observations} \\tablehead{ \\colhead{Transition} & \\colhead{$\\nu_{line}$} & \\colhead{$E_u$} & \\colhead{Data} & \\multicolumn{2}{c}{Synthesized Beam} & \\colhead{RMS}  & \\colhead{Figures}\\\\ \\colhead{} & \\colhead{(GHz)} & \\colhead{(K)} & \\colhead{} & \\colhead{$\\arcsec\\times\\arcsec$} & \\colhead{$\\degr$} & \\colhead{(Jy~beam$^{-1}$)} & \\colhead{} \\\\ } \\startdata \\lsbii   & 214.689380 & 147.9 &  LR & $8\\farcs3\\times3\\farcs1$ & 34.2 & 0.065 & \\ref{Fspec_so2} \\\\ &            &       & ROB & $1\\farcs3\\times0\\farcs9$ & 59 & 0.085 & \\ref{Fcubos} \\\\ &            &       &  HR & $0\\farcs7\\times0\\farcs4$ & 13.2 & 0.085 & \\ref{Fmoms} \\\\ \\lsbiii  & 214.728330 & 229.1 &  LR & $8\\farcs3\\times3\\farcs1$ & 34.2 & 0.065 & \\ref{Fspec_so2} \\\\ \\usbv    & 223.883569 &  58.6 &  LR & $7\\farcs9\\times3\\farcs0$ & 34.2 & 0.065 & \\ref{Fspec_so2} \\\\ \\usbiv   & 224.264811 & 207.9 &  LR & $7\\farcs9\\times3\\farcs0$ & 34.2 & 0.065 & \\ref{Fspec_so2} \\\\ \\usbii   & 225.153702 &  93.1 &  LR & $7\\farcs9\\times2\\farcs9$ & 34.2 & 0.065 & \\ref{Fspec_so2} \\\\ &            &       & ROB & $1\\farcs2\\times0\\farcs9$ & 58 & 0.070 & \\ref{Fcubos} \\\\ &            &       &  HR & $0\\farcs7\\times0\\farcs4$ & 12.9 & 0.080 & \\ref{Fmoms} \\\\ \\lsbi    & 214.357004 &  81.3 &  LR & $8\\farcs3\\times3\\farcs1$ & 34.2 & 0.065 & \\ref{Fspec_so2} \\\\ \\lsbiv   & 215.220653 &  44.1 &  LR & $8\\farcs2\\times3\\farcs1$ & 34.2 & 0.070 & \\ref{Fspec_so2} \\\\ &            &       & ROB & $1\\farcs2\\times0\\farcs9$ & 58 & 0.080 & \\ref{Fcubos} \\\\ &            &       & TAP & $4\\farcs0\\times2\\farcs2$ & 36.7 & 0.080 & \\ref{Fsocube},\\ref{Fmomso},\\ref{Fespec} \\\\ \\co      & 224.714385 &  16.2 &  LR & $7\\farcs9\\times2\\farcs9$ & 34.2 & 0.075 & \\tablenotemark{(a)} \\\\ &            &       & TAP & $4\\farcs0\\times2\\farcs1$ & 37.8 & 0.080 & \\ref{Fmomusb3} \\\\ \\formal  & 225.697775 &  33.5 &  LR & $7\\farcs9\\times2\\farcs9$ & 34.2 & 0.090 & \\ref{Fspec_so2} \\\\ &            &       & ROB & $1\\farcs2\\times0\\farcs9$ & 58 & 0.100 & \\ref{Fcubos} \\\\ &            &       & TAP & $4\\farcs0\\times2\\farcs1$ & 37.2 & 0.095 & \\ref{Fformalcube},\\ref{Fmomformal},\\ref{Fespec} \\\\ \\enddata \\tablecomments{Data extracted from JPL catalogue. The fourth column \\textit{Data} is a code giving the resolution of each map. LR means low angular resolution data, HR, high angular resolution data, ROB, combined images with the robust set to 0.3 and TAP, combined images with a taper applied and restricted in the uvplane (see text, section \\S2).} \\tablenotetext{(a)}{The compact configuration data of the \\co is not shown in any figure. However, it is used to obtain the velocity, density and mass of the dense core of \\ggd (see section \\S 3.1).} \\label{Tsummary} \\end{deluxetable*}  \\begin{figure} \\epsscale{1} \\plotone{fig1} \\caption{Contour image of the zero order moment superposed on the first order moment colour image of the \\co line emission (combined compact and extended configuration data with a Gaussian uv taper) toward \\ggd. Contour levels are -9, -4, 4, 9, 15, 30, 45, 60, 75 and 90 times  0.10~Jy~beam$^{-1}$\\kms, the rms of the image. The units of the scale bar are in \\kms. The molecular emission is integrated over velocities between 9 and 18\\kms. The triangles mark the position of the millimeter sources MM1 and MM2, and the position of the molecular core, MC. The dashed arrows show the directions of the NW and SE outflows which seem to originate at the MM2 position and also the NE-SW thermal radio jet, launched from MM1. The synthesized beamsize is shown at the bottom right corner and the primary beam is indicated by the dashed circle. } \\label{Fmomusb3} \\end{figure}  \\begin{figure} \\epsscale{1} \\plotone{fig2} \\caption{Contour image of the zero order moment of the \\lsbiv line emission (combined compact and extended configuration data with a Gaussian uv taper) superimposed on a K--band near infrared image of the reflection nebula in the \\ggd region (courtesy of Thomas Geballe; see \\citealt*{1992Aspin}). The infrared polarimetric study carried out by \\cite{1987Yamashita} demonstrated that the scattered light from the infrared reflection nebula occurs on grains at the walls of a parabolic cavity rather than on grains inside the whole outflow lobe. Contour levels are -9, -4, 4, 9, 15, 30, 45, 60, 75 and 90 times 0.17~Jy~beam$^{-1}$\\kms, the rms of the image. Symbols are as in Fig. \\ref{Fmomusb3}. } \\label{Fmomso} \\end{figure}  \\begin{figure*} \\epsscale{1} \\plotone{fig3} \\caption{Channel velocity image of the \\lsbiv extended emission (combined compact and extended configuration data with a Gaussian uv taper) from the central region of \\ggd. The emission is averaged in velocity bins of 1.14\\kms. The velocity of each channel is indicated in the top left corner of each panel in \\kms units. The cloud velocity is at about 11.8\\kms. The contour levels are -6, -4, 4, 6, 9, 13, 18, 24, 31, 39, 48 and 58 times 0.08~mJy~beam$^{-1}$, the rms of the image. Symbols are as in Fig. \\ref{Fmomusb3}. The synthesized beam size is shown in the bottom left-hand corner of the most redshifted channel. The blueshifted emission appears to trace mainly the walls of the cavity excavated by the radio jet from MM1. At the position of MM1 the image shows emission probably coming from the rotating disk/ring, absorbed in the bluest channels by the extended emission. The redshifted channels show also compact emission toward the position of MC. } \\label{Fsocube} \\end{figure*}  \\begin{figure} \\epsscale{0.7} \\plotone{fig4} \\caption{\\lsbiv (top panel) and \\formal (bottom panel) spectra toward the center position of MC (black line). Spectra toward the clump located to the north--west of MC (green line) are included to compare with that of MC. The dotted line marks the systemic velocity (11.8\\kms). The intensity is in Kelvin. The spectra are extracted from the combined compact and extended configuration data with a Gaussian uv taper. } \\label{Fespec} \\end{figure}  \\begin{figure} \\epsscale{1} \\plotone{fig5} \\caption{Contour image of the zero order moment of the \\formal line emission (combined compact and extended configuration data with a Gaussian uv taper) overlaid with a K--band near infrared image of the \\ggd region (as shown in Fig. \\ref{Fmomso}). Contour levels are -9, -4, 4, 9, 15, 30, 45, 60, 75 and 90 times 0.33~Jy~beam$^{-1}$\\kms, the rms of the image. Symbols are as in Fig. \\ref{Fmomusb3}. } \\label{Fmomformal} \\end{figure}  \\begin{figure*} \\epsscale{1} \\plotone{fig6} \\caption{Channel velocity image of the \\formal extended emission (combined compact and extended configuration data with a Gaussian uv taper). The emission is averaged in velocity bins of 1.14\\kms. The velocity of each channel is indicated at the top left corner of each panel in \\kms units. The cloud velocity is at about 11.8\\kms. The contour levels are -6, -4, 4, 6, 9, 13, 18, 24, 31, 39, 48 and 58 times 0.095~mJy~beam$^{-1}$, the rms of the image. Symbols are as in Fig. \\ref{Fmomusb3}. } \\label{Fformalcube} \\end{figure*}  In recent years, several works have provided evidence for the presence of collimated jets (HH~80-81, \\citealt{1993Marti}; Cep~A, \\citealt{2006Curiel}; IRAS~23139, \\citealt{2006Trinidad}; IRAS~16547, \\citealt{2008Rodriguez}), molecular outflows (see a summary in \\citealt{2005Zhang,2007Arce}), and flattened  structures of dust and gas (some with rotation and some with infall signposts;  e.g., Cep~A, \\citealt{2005Patel}; IRAS~20126, \\citealt{2005Cesaroni}; AFGL~2591, \\citealt{2006VanderTak}; W51 North, \\citealt{2008Zapata};  IRAS~16547, \\citealt{2009Franco-Hernandez}; W33A-MM1, \\citealt{2010GalvanMadrid}), surrounding high-mass protostars.  The evidence leads to the interpretation that massive star formation is analogous  to low-mass star formation, that is, via accretion from a flat rotating disk, with a jet of ionized material, and with an associated molecular outflow. However, the physical  characteristics of the possible accretion disks and their associated outflows  have not been well characterized yet. Hence, the formation process of massive stars  remains an open issue. There are several problems that hamper the analysis of massive star formation regions (MSFRs).  First, massive protostars form in clusters, which together with projection effects, make these regions difficult to interpret. Moreover, the  MSFRs are found at larger distances than low-mass star formation regions, with  typical distances of 2-5~kpc. Thus, to get an insight of the closest regions to the protostars, where accretion disks are expected, the highest angular resolution of the current  telescopes is needed. Even so, most of the dust and gas structures detected around high-mass protostars are $\\sim1000$~AU. This is the smallest diameter that an interferometer can resolve with a 1$\\arcsec$ angular resolution at those distances. At these scales, such structures could harbour systems of several protostars, that could be forming high mass stars by other mechanisms such as mergers. In fact, the nearest massive protostars (Orion BN/KL and Cepheus A HW2), show complex scenarios, with a possible merger of protostars (Orion BN/KL, \\citealt{2009Zapata}), or close passages between protostars (Cep~A HW2, \\citealt{2009Cunningham}). Therefore, with the current interferometers, only a few MSFRs may be studied with adequately high angular resolution.  To demonstrate the existence of a disk around a massive protostar, it is not enough to detect its (sub)millimeter dust continuum emission;  its kinematics must also be measured via the molecular  line emission associated with the gas--phase chemistry developed  after the evaporation of the molecules from the ice mantles  (\\citealt{1998Hatchell,2007Cesaroni}). However, searching for molecular tracers of disks in MSFRs is a complex task. There are molecules that show emission from the envelope and the disk simultaneously. Other molecules show optically thick emission, thus complicating the kinematic study of the disk.  On the other hand, S-bearing species (such as H$_2$S, SO, SO$_2$, CS, OCS...) could be intimately linked with the evaporation process of the disk surface (\\citealt{1997Charnley,1998Hatchell,2003VanderTak,2005MartinPintado}), becoming  good tracers of the dynamics of the innermost parts of the high mass protostars.  Recently, several papers have been published on the detection of S-bearing species in disks and other warm gas-structures of MSFRs (\\citealt{2006VanderTak,2007JimenezSerra,2009Klaassen,2009Franco-Hernandez, 2009Zapata}). In particular, SO$_2$ transitions, ubiquitous within the (sub)millimeter range, show a very compact nature, suggesting their close association with circumstellar structures.  The GGD27 complex is an active star forming region located at a distance of 1.7~kpc. It shows a spectacular (5.3~pc long)  and highly collimated thermal radio jet (\\citealt{1993Marti,1995Marti,1998Marti}), embedded within a powerful CO bipolar outflow (\\citealt{1989Yamashita,2004Benedettini}). The jet has a position angle of $\\sim21\\degr$ (\\citealt{1993Marti,1999Marti}) and ends at two southern and very bright Herbig-Haro objects (HH~80 and 81, originally discovered by \\citealt{1988Reipurth}) and at a radio source to the north (HH~80 North,  \\citealt{1994Girart}). Linear polarization has been detected in the radio emission  from this jet, indicating the presence of a magnetic field coming from the disk (\\citealt{2010Carrasco}).The central part of the radio jet has a bright  far-infrared counterpart (\\ggd), that implies the presence of a luminous young star ($\\sim2\\times10^4$~L$_{\\sun}$) or a cluster of stars (\\citealt{1992Aspin,1997Stecklum}).  Submillimeter and millimeter wavelength observations of the central part of the thermal radio jet have revealed two dusty sources, MM1 and MM2, separated by about 7$\\arcsec$ (\\citealt{2003Gomez, 2009Qiu}), which are apparently in very different evolutionary stages (\\citealt{2011Fernandez}, hereafter Paper I). MM1, the south-western source, coincides with the origin of the thermal radio jet and has a dust temperature of 109~K. It has a weak extended envelope, possibly surrounding a very compact (R$\\lesssim$150~AU) disk-protostar system. The mass of the disk derived from the continuum emission is $\\sim4.1$~\\msun. Such a massive disk could have an extremely high accretion rate ($10^{-3}-10^{-2}$\\msun~yr$^{-1}$). However, the case for an accretion disk in MM1 has not been unambiguously confirmed yet. The bolometric luminosity obtained from fitting the SED of high angular resolution data is $\\geq3300$\\lsun, similar to that of a B1 Zero Age Main Sequence (ZAMS) star. However, the massive disk of MM1 (compared to those of low-mass protostars), the presence of the powerful outflow and the youth of the protostar (it has not developed a compact H~II region yet), indicate that MM1 could become a B0 or a more massive star. MM1 is also associated with compact emission from several high density tracers, such as CS (\\citealt{1991Yamashita}), SO (\\citealt{2003Gomez}), CH$_3$OH, CH$_3$CN, H$_2^{13}$CO, OCS, HNCO and SO$_2$ (\\citealt{2009Qiu}).  The physical characteristics of MM2 are typical of the Class 0 low--mass protostars (see e.g., \\citealt{1993Andre}), except for its total  mass of at least 11~M$_{\\sun}$ (4-5 times more mass than a typical low-mass Class 0 source, see e.g. \\citealt{2006Girart,2009Rao}). MM2 appears as an extended source  when observed with low angular resolution at 1.4~mm, but it is separated into two compact sources when observed with high angular resolution (Paper I). The stronger component in MM2 spatially coincides with very weak free-free emission, suggesting that it could be a high/intermediate mass protostar. The weaker component in MM2 could be a pre-protostellar core. Previous CO observations (\\citealt{2009Qiu}) have associated the position of MM2 with the origin of a young outflow running to the east. In addition, another possible CO outflow, running north-west of MM2 could also have originated close to this position.  There is also a possible molecular core (hereafter MC, \\citealt{2009Qiu}) located about $4\\arcsec$ to the north-west of MM2. Although MC is traced by several molecules (CH$_3$CN, H$_2^{13}$CO, OCS y HNCO, \\citealt{2009Qiu}, and a Class I CH$_3$OH maser, \\citealt{2004Kurtz}), it is not associated with a bright and compact millimeter source (Paper I). The diverse line emission from MC has been interpreted by \\cite*{2009Qiu} as the emission from a hot core warmed by the radiation of a young and massive protostar.  Here we present an analysis of the molecular emission of several lines (mostly S-bearing species) detected in the central region of \\ggd using 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; \\citealt{2004Ho}). The SMA allowed us to detect several transitions simultaneously, some of which seem to trace a circumstellar disk/ring around MM1, with subarcsecond resolution  ($\\sim0\\farcs5$), permitting a beam size of $\\sim850$~AU at the source distance.  This resolution is necessary in order to perform a reliable kinematical analysis of the motions of the circumstellar disks at distances greater than 1--2~kpc.  The continuum emission of the observations was presented in a recent publication (Paper I). In section 2 we briefly describe the observations, while in section 3 the results of the observations are given. Section 4 contains the analysis of the column density and the  temperature of the disk/ring of MM1, followed by a discussion in section 5. Finally, in section 6 we draw our main conclusions. ",
        "conclusions": "We have performed an SMA observational study of the molecular gas toward the central region of \\ggd, and have obtained the following main results.  \\begin{itemize} \\item The \\lsbii and \\usbii data show a compact (R$\\simeq425$--850~AU) molecular disk--like structure spatially coincident with the position of the compact millimeter dusty  source MM1 and the radio jet located at the \\ggd position. The molecular structure is perpendicular to the radio jet axis and it shows a clear velocity gradient that we interpret as rotation. The dynamical mass (11--15\\msun) allowed us to derive the mass of the protostar (7--11\\msun), making use of the  mass that was inferred from the continuum data (Paper I).  The RADEX analysis constrains the disk--like structure temperature between 120 and 160~K  and its volume density as $\\gtrsim2\\times10^9$\\cmt.  We have also found that the rotating system could be unstable due to  gravitational disturbances.  \\item The \\co emission shows an elongated structure of about 0.2~pc~$\\times$~0.1~pc  with an orientation roughly perpendicular to the radio jet. This emission  seems to belong to the dense core in which the protostars associated with the MM1, MM2 and MC sources are forming.  \\item The H$_2$CO and the SO emission appears clumpy and is mainly enhanced toward  the position of MM1 and MC. It is possibly tracing in part the cavity walls  excavated by the thermal radio jet. At the position of MC, it shows a line profile with a red wing extending up to tens of \\kms. This line profile is probably due to shocked gas entrained by the outflows in the region. We speculate that a possible scenario is the interaction between the outflow associated with MM2 and the dense molecular core at this position. \\end{itemize}"
    },
    "1107/1107.3481_arXiv.txt": {
        "abstract": "The Maximum Entropy Method (MEM) for the deconvolution of radio interferometry images is mathematically well based and presents a number of advantages over the usual CLEAN deconvolution, such as appreciably higher resolution. The application of MEM for polarisation imaging remains relatively little studied. CLEAN and MEM intensity and polarisation techniques are discussed in application to recently obtained 18cm VLBA polarisation data for a sample of Active Galactic Nuclei. ",
        "introduction": "Consider a radio source which gives rise to a brightness distribution, $I(x,y)$, in the radio sky of the Earth. There exists a Fourier transform (FT) relationship between the true brightness distribution and the complex visibility function $V(u,v)$ measured by a VLBI array as a function of the baseline vector.  \\begin{equation} V(u,v)=F.T.(I(x,y)) \\end{equation} \\noindent It is not a simple case of taking the inverse Fourier transform of the visbilities to return to the true brightness distribution as, due to the limited uv coverage offered by Earth rotation synthesis VLBI, most components of the visibility function are not measured. Mathematically this corresponds to multiplying the total visibility by a sampling function, $S(u,v)$, which eliminates most of the visibilities and leaves the observer with only $V_{sampled}(u,v)$. This means that, even before noise is taken into account, any attempt to recover the sky brightness distribution from the measured uv visibilites is a mathematically ill-posed problem.  \\begin{equation} V_{sampled}(u,v)= V(u,v) S(u,v) + Noise \\end{equation} \\noindent Deconvolution algorithms are used to attempt to interpolate the information contained in the measured visibilities to reconstruct the unmeasured visibilities and get as true an image of the sky brightness distribution as possible. There is no single ``best'' deconvolution algorithm, however the two most common are outlined below.    \\subsection{The CLEAN Algorithm}  The CLEAN algorithm (H\\\"ogbom, 1974) is a standard technique used to deconvolve the measured visibilities and attempt to recover the data from the unmeasured visibilities. First the measured visibilities are Fourier transformed to create a ``dirty'' map (containing only the information from the measured visibilities). This map is then described by a set of ``CLEAN'' components ($\\delta$ functions). The set of ``CLEAN'' components is then convolved with the CLEAN beam (a Gaussian fit to the primary lobe of the point spread function) and the residual noise added in to give the final ``clean'' image. Images of Stokes I,Q and U parameters can be made independently using this technique.  \\subsection{The Maximum Entropy Method (MEM)}  A second way to deconvolve the visibilities and recover some of the unmeasured data is to make a model intensity map of the source that satisfies specified mathematical criteria. This model can then be varied until the simulated visibilities agree with the actual data to within a specified limit, thus maximising the similarity of the model intensity and the true map. However, as the model is changed one cannot let the model visibilities converge to the actual measured visibilities, or the model would then simply yield the original ``dirty'' map. A regularising parameter is required to stop this convergence. In MEM, this parameter is the entropy of the image. Consider the following function:  \\begin{equation} J=H(I_{m},P_{m})-\\alpha \\chi_{I}^{2}(V_{m},V_{d})-\\beta \\chi_{P}^{2}(V_{m},V_{d}) - conditions \\end{equation} \\noindent where H is the entropy of the model map of the source and $\\chi^{2}$ is a measure of the difference between the model (subscript m) and observed (subscript d) visibilities (there are two such terms, one for intensity, $\\chi_{I}^{2}(V_{m},V_{d})$, and a second for polarisation, $\\chi_{P}^{2}(V_{m},V_{d})$. $\\alpha$ and $\\beta$ are Lagrange parameters, and other conditions are also included which represent additional constraints, such as positivity of the Stokes I component. The optimal model of the source maximises the function J above. This results in a balance between entropy (representing noise, and the effect of unsampled visibilities) and fidelity to the observed data.\\\\  Multiple forms of entropy can be used; common choices include Shannon entropy (\\ref{shannon}), which is suitable for unpolarised emission, and the Gull and Skilling entropy (\\ref{gull}), a generalised form of the Shannon entropy, suitable for polarised emission.  \\begin{equation} \\label{shannon} H=-\\sum_{k}I_{k} log(\\frac{I_{k}}{B_{k}}) \\end{equation}  \\begin{equation} \\label{gull} H=-\\sum_{k}I_{k} (log(\\frac{I_{k}}{e B_{k}})+\\frac{1+m_{i}}{2}log(\\frac{1+m_{i}}{2})+\\frac{1-m_{i}}{2}log(\\frac{1-m_{i}}{2})) \\end{equation} \\noindent In these equations, the index $k$ represents summation over pixels, and $I_{k}$ and $m_{k}$ are the intensity and fractional polarisation, respectively, at pixel $k$. $B$ is a bias map, to which the solution defaults in the absence of data. See Holdaway \\cite{holdaway} and Gull and Skilling \\cite{gas} for more details about the MEM and different forms of polarisation entropy. ",
        "conclusions": "Both CLEAN and MEM offer unique and complementery perspectives on the same uv visibility data. MEM appears less sensitive to low intensity emission, but provides higher resolution, allowing jet direction and morphology to be studied in more detail. There is good agreement between the CLEAN and approximate MEM polarisation maps we have constructed, suggesting that the approximations made in the work-around used to produce polarised MEM maps in AIPS have not adversely affected the polarisation. However, in order to generate a truely reliable MEM polarisation map a form of polarised entropy must be used and we are in the process of implementing the entropy (\\ref{gull}) for this purpose. The synergy between CLEAN and MEM in studying polarised emission should then lead to a clearer picture of the polarisation intensity and direction along the jet. We also plan to investigate the construction of MEM spectral-index and Faraday-rotation maps based on multi-frequency intensity and polarisation VLBI maps."
    },
    "1107/1107.1484_arXiv.txt": {
        "abstract": "We report new observations of the unusually active, high proper motion L5e dwarf 2MASS J13153094$-$2649513. Optical spectroscopy with Magellan/MagE reveals persistent nonthermal emission, with narrow H~I Balmer, Na~I and K~I lines all observed in emission. Low-resolution near-infrared spectroscopy with IRTF/SpeX indicates the presence of a low-temperature companion, which is resolved through multi-epoch laser guide star adaptive optics imaging at Keck.  The comoving companion is separated by 338$\\pm$4~mas, and its relative brightness ($\\Delta{K_s}$ = 5.09$\\pm$0.10) makes this system the second most extreme flux ratio very low-mass binary identified to date. Resolved near-infrared spectroscopy with Keck/OSIRIS identifies this companion as a T7 dwarf. The absence of Li~I absorption in combined-light optical spectroscopy constrains the system age to $\\gtrsim$0.8--1.0~Gyr, while the system's kinematics and unusually low mass ratio (M$_2$/M$_1$ = 0.3--0.6) suggests that it is even older. A coevality test of the components also indicates an older age, but reveals discrepancies between evolutionary and atmosphere model fits of the secondary which are likely attributable to poor reproduction of its near-infrared spectrum. With a projected separation of 6.6$\\pm$0.9~AU, the {\\namesh} system is too widely separated for mass exchange or magnetospheric interactions to be powering its persistent nonthermal emission.  Rather, the emission is probably chromospheric in nature, signaling an inversion in the age-activity relation in which strong magnetic fields are maintained by relatively old and massive ultracool dwarfs. ",
        "introduction": "Nonthermal emission is commonly observed among the lowest-mass stars, traced by optical line (e.g., Ca II, H I sequence),  X-ray, UV and radio emission.  This emission can be both persistent (quiescent) and eruptive, with short-duration flares from M dwarfs occurring at a rate of roughly 3\\%  \\citep{2010AJ....140.1402H}.  The incidence and strength of quiescent magnetic activity as traced by the H$\\alpha$ line reaches $\\gtrsim$80\\% and {\\lhalbol} $\\approx$ $-$4, respectively, among nearby late-type M dwarfs \\citep{2000AJ....120.1085G,2004AJ....128..426W,2011AJ....141...97W, 2007AJ....133.2258S}, but both metrics decline precipitously for the cooler L and T dwarfs \\citep{2002AJ....123.2744B, 2007AJ....133.2258S, 2010AJ....139.1808S} and sources far from the Galactic plane \\citep{2006AJ....132.2507W, 2008AJ....135..785W}. Similar declines are seen in X-ray emission, but surprisingly not at radio frequencies \\citep{2002ApJ...572..503B, 2005ApJ...626..486B, 2006ApJ...648..629B, 2008A&A...487..317A}. Assuming that this nonthermal emission arises from magnetic interaction, the decline with vertical scale height among M dwarfs is likely an age effect, as angular momentum loss over time results in weakened magnetic dynamos.  However, spin-down timescales exceed 5~Gyr beyond spectral type M5 \\citep{2008AJ....135..785W}, and many late-type dwarfs are found to be rapid rotators (P $\\approx$ 2-10~hr; \\citealt{2003ApJ...583..451M, 2006ApJ...647.1405Z, 2008ApJ...684.1390R}) exhibiting kilogauss magnetic fields at their photospheres \\citep{2007ApJ...656.1121R, 2008ApJ...684..644H}. Hence, the decline in magnetic emission with spectral type must arise from a different effect.  The favored cause is the decoupling of cooler, increasingly neutral photospheres from magnetic structures, which reduces magnetic stresses and the frequency of magnetic reconnection above the (sub)stellar surface (e.g., \\citealt{1999A&A...341L..23M, 2002ApJ...571..469M, 2002ApJ...577..433G}). Deeper reconnection events may continue to power strong flaring bursts observed in a handful of weakly active or inactive late-M and L dwarfs (e.g., \\citealt{1999ApJ...527L.105R, 2003AJ....125..343L, 2007AJ....133.2258S}).  Contrary to these trends, a very rare set of low-temperature ``hyperactive'' dwarfs exhibit unusually prodigious and persistent nonthermal emission. One of the first examples of these to be identified was {\\name} (hereafter {\\namesh}; \\citealt{2002ApJ...564L..89H, 2002ApJ...575..484G}), a high proper motion L5e dwarf which has exhibited sustained but variable H$\\alpha$ emission on no fewer than seven epochs spanning nearly a decade \\citep{2002ApJ...564L..89H, 2002ApJ...580L..77H, 2002ApJ...575..484G, 2003AJ....125..343L, 2005A&A...439.1137F, 2006A&A...459..511B, 2008ApJ...689.1295K}. With {\\lhalbol} $\\approx$ $-$4, {\\namesh} is as active as a mid-type M dwarf, but is 1--2 orders of magnitude brighter in H$\\alpha$ than comparably-classified L dwarfs \\citep{2000AJ....120.1085G, 2007AJ....133.2258S}. In addition, H$\\beta$ and Na~I D lines have also been observed in emission \\citep{2005A&A...439.1137F}. Since the photosphere of {\\namesh} is cool and likely to be highly neutral, the origin of its unexpected emission remains a mystery.  Neither the kinematics nor spectral characteristics of {\\namesh} indicate youth, and the absence of mid-infrared excess argues against accretion from a protoplanetary or debris disk \\citep{2007ApJ...659..675R}. Other hyperactive dwarfs, such as the L1e dwarf 2MASS~J10224821+5825453 ({\\lhalbol} $\\approx$ $-$3.5 to $-$2.7; \\citealt{2007AJ....133.2258S}) and the T6.5e dwarf 2MASS J1237392+652615 ({\\lhalbol} $\\approx$ $-$4.6 to $-$4.2; \\citealt{2000AJ....120..473B, 2002AJ....123.2744B, 2007ApJ...655..522L}; hereafter 2MASS~J1237+6526) also lack evidence of disk accretion and do not appear to be particularly young.  Alternative mechanisms, such as acoustic heating, unusually strong magnetic fields, and Roche lobe overflow from a substellar companion have been proposed, but none of these scenarios have been validated.  In this article, we report new observations of {\\namesh} that reveal both continued nonthermal optical line emission in several neutral atomic species, and the presence of a T dwarf companion at a projected separation of 7~AU. In Section~2 we describe our combined-light optical and near-infrared spectroscopic observations, the  latter of which yields preliminary evidence for  a brown dwarf companion. In Section~3 we describe adaptive optics (AO) imaging and spectroscopic observations that confirm the presence of the companion and allow determination of its separation and classification. In Section~4 we analyze the observed and inferred physical properties of the components, the latter based on comparison to evolutionary and atmospheric models. We also perform a coevality test to examine the reliability of these models. In Section~5 we discuss how the properties of {\\namesh} argue for a magnetic origin to its persistent emission, powered by a strong magnetic field retained by a relatively old and massive cool dwarf. We summarize our results in Section~6. ",
        "conclusions": "We have identified a T7 brown dwarf companion to the unusually active L5e dwarf {\\namesh}, a source that continues to exhibit strong H~I and alkali line emission despite its late spectral type. Resolved imaging and spectroscopy confirm the pair to be co-moving and co-spatial, and evolutionary models indicate an unusually low mass ratio as compared to other low-mass field binaries. The spectral and kinematic properties of both components confirm prior indications that this system is relatively old ($\\tau$ $\\gtrsim$ 0.8--1.0~Gyr), and likely a member of the old Galactic disk population. Joint comparison to atmospheric and evolutionary models (a coevality test) also supports an older age for this system, but reveals discrepancies in the case of the secondary; we suspect these are due to continued shortcomings in the modeling cool brown dwarf spectra. The age and separation of the system, coupled with the narrow emission lines and absence of mid-infrared excess, rule out accretion from or magnetic interaction with the T dwarf secondary as the emission source. Rather, we attribute it to an unusually strong magnetic field as predicted by energy flux scaling arguments for relatively old and massive low-mass dwarfs. A direct test of this hypothesis can be achieved through Zeeman line broadening measurements, although the source of chromospheric heating remains a separate issue.  Whether the larger sample of hyperactive cool dwarfs are also old, relatively massive, and possess strong magnetic fields remains to be determined, but they hint at a remarkable inversion of the standard age-activity relationship for low-mass stars."
    },
    "1107/1107.4742_arXiv.txt": {
        "abstract": "We report the discovery of a unique radio galaxy at $z$=0.137, which could possibly be the second spiral-host large radio galaxy and also the second triple-double episodic radio galaxy. The host galaxy shows signs of recent star formation in the UV but is optically red and is the brightest galaxy of a possible cluster. The outer relic radio lobes of this galaxy, separated by $\\sim$1 Mpc, show evidence of spectral flattening and a high fraction of linear polarisation.  We interpret that these relic lobes have experienced re-acceleration of particles and compression of the magnetic field due to shocks in the cluster outskirts.  From the morphology of the relics and galaxy distribution, we argue that re-acceleration is unlikely to be due to a cluster-cluster merger and speculate about the possibility of accretion shocks.  The source was identified from SDSS, GALEX, NVSS and FIRST survey data but we also present follow up optical observations with the Lulin telescope and 325 MHz low frequency radio observations with the GMRT. We briefly discuss the scientific potential of this example in understanding the evolution of galaxies and clusters by accretion, mergers, star formation, and AGN feedback. ",
        "introduction": "\\begin{figure*} \\hbox{ \\psfig{file=SPECA_LULIN_GALEXNUV_GALEXFUV-abc-gimp.eps,width=4.1in,angle=-90} \\psfig{file=SPECA_Lulin_GALFIT_results_residual_single-d.ps,width=2.6in,angle=0} } \\caption[]{(a) Lulin $R$-band image of the host galaxy Speca, presented both in colour (linear scale) and contours (logarithmic scale). (b) The same $R$-band contours are superposed on the smoothed near-UV image from GALEX in linear scale. (c) The same $R$-band contours are superposed on the smoothed far-UV image also from GALEX in linear scale. (d) Radial surface brightness profile of the Lulin $R$-band image. Black and red data are shown for the SW and NE parts of the disk, respectively, for comparison. Fitted model profiles are of a nuclear point source and exponential disk (cyan continuous line) and of a de Vaucouleurs profile (green broken line). For the former model, both individual model components and the total are shown. } \\end{figure*} By incorporating processes of AGN feedback into the models of hierarchical structure formation, it is possible to explain observed galaxy population properties, such as the galaxy and cluster luminosity function, the $M_\\bullet$$-\\sigma$ relation, and the colour-magnitude bi-modality of nearby galaxies (Springel, Di Matteo \\& Hernquist 2005, Croton et al. 2006).  Powerful and large-scale feedback in the form of radio lobes seen as multiple cavities in X-ray images of clusters has provided strong support for such models (Nulsen et al. 2005, Sanders \\& Fabian 2007, Randall et al. 2011).  Luminous radio galaxies are also found to be tracers of high-$z$ cluster formation (Venemans et al. 2007, Zirm et al. 2005).  In the nearby Universe, diffuse Mpc-scale shell-like radio emission found in the outskirts of clusters has been interpreted as emission from acceleration of particles at shock fronts due to cluster mergers or structure formation (Ensslin et al. 1998; Ensslin et al. 2001; Bagchi et al. 2006, 2011; van Weeren et al. 2010).  Apart from mergers, clusters grow by accretion from cosmic galaxy filaments. Evidence of shocks, predicted for such accretion, is yet to be convincingly detected from X-ray observations, but many radio observations show promising prospects (Ensslin et al. 1998, Pfrommer, Ensslin \\& Springel 2008).  A wide-angle tail radio galaxy recently discovered to be present in a several-Mpc-scale galaxy filament has been proposed as evidence of gas accretion through the filament onto the cluster (Edwards et al. 2010).  Due to the Mpc-scale size and episodic nature of radio galaxies (Saikia \\& Jamrozy 2009), whose relic plasma as old as 2 Gyr can be revived (Ensslin \\& Gopal-Krishna 2001), low frequency radio observations, naturally tracing old relativistic plasma, are powerful probes of these merger, accretion, and feedback processes driving cluster formation and evolution.  The time-scale of 1$-$2 Gyr is comparable to the dynamical timescale of galaxies in the clusters, the timescale of galaxy mergers, the timescale of the star formation in UV-bright early-type galaxies (Yi et al. 2005) and quenching of star formation in the post-starburst (E+A) galaxies (Goto 2005).  Therefore, combined low-frequency study of relics of AGN feedback and recent star formation in AGN host galaxies is a potential new strategy of investigating galaxy-black hole co-evolution via mergers and feedback.  Here we report the discovery of a unique radio galaxy whose host galaxy appears to be a spiral or disk with young star formation, rather than a typical giant elliptical galaxy with a predominantly old stellar population. Equally unique is that it shows signs of three distinct episodes of jet ejection where the outer relic radio lobes, separated by $\\sim$1 Mpc, have experienced particle re-acceleration possibly due to shocks in the cluster outskirts. The optical source had been cataloged before as SDSS J140948.85-030232.5 or LCRS B140713.3-024824 or MaxBCG J212.45357-03.04237 (being brightest in the cluster) but for ease in referring we nicknamed it, due to its uniqueness, `Speca' (SPiral-host Episodic Cluster-dominant AGN). For the spectroscopic redshift of Speca $z$=0.1378, 1 arcsec corresponds to  2.396 kpc, in a Universe with $H_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}$=0.27 and $\\Omega_{\\rm vac}$=0.73. ",
        "conclusions": "\\begin{figure} \\vbox{ \\psfig{file=SPECA-GMRT-ACTIVE-SE-RELIC.PS,width=2.5in,angle=0} \\psfig{file=Speca-combined-colour.eps,width=2.5in,angle=0} } \\caption[]{ Top Panel: The high resolution ($\\sim$9 arcsec) GMRT 325 MHz image, both in contour and grey-scale, of the inner, middle and SE outer relic lobes of Speca. The isophotal major-, minor-axis and P.A of the host galaxy, as in the text, are also indicated with an ellipse and '+' sign.  Bottom Panel: The low resolution (45 arcsec) GMRT 325 MHz image in contours is superposed on the $R$-band optical image. The relic lobes and an unrelated radio source are marked. The spectroscopically confirmed members of the cluster are marked by boxes, and possible members from SDSS $z_{\\rm photo}$ are marked by triangles. The inset shows details of the NW outer relic lobe with NVSS data shown as blue contours, high resolution ($\\sim$9 arcsec) GMRT 325 MHz image shown in red contours, and the low resolution (45 arcsec) 325 MHz image as in the background is shown in thick black contours.} \\end{figure} \\label{optical} \\subsection{Disk structure in the host galaxy} While comparing Sloan Digital Sky Survey (SDSS) and Faint Images of the Radio Sky at Twenty-Centimeters (FIRST) images, we identified the host galaxy Speca as an unusual object. The galaxy lies at the geometrical centre of the radio lobes, and the SDSS spectrum from the central region shows a [N{\\small II}]/H$\\alpha$ line ratio higher than unity, suggesting it to be the AGN host galaxy. Our $R$-band image shows that the main body clearly suggests a morphology consistent with a nearly edge-on disk galaxy (Figure 1 (a)). The north-east (NE) side of the nucleus shows a curved feature similar to a spiral-arm seen nearly edge-on.  Similar structure is also seen in SDSS $g'$, $r'$ and $i'$-band images, but any possible counterpart of this spiral arm on the south-western side of the nucleus is not seen. The curved nature of the spiral arm-like feature and extended isophotal contours on the southern side, disfavours any mid-plane dust causing the observed stellar structure. Measured up to 25 mag arcsec$^{-2}$ in the SDSS $r'$-band, the size of the galaxy is 60 $\\times$ 29 kpc (25 arcsec and 12 arcsec with P.A. 21$^{\\circ}$) suggesting that Speca is a large, mature host galaxy.  In Figure 1 (b, c) we also present a comparison of the $R$-band image with near-UV (NUV) and far-UV (FUV) images available on-line from the Galaxy Evolution Explorer (GALEX). Surprisingly, the UV emission (smoothed with a tophat kernel of radius 2 (4) pixels for NUV (FUV)) correlates nicely with the optical major-axis. The lack of dominant UV emission from the centre of the galaxy suggests that the emission is associated with star formation and unlikely to be dominated by the central AGN.  In the NUV the northern side is brighter than the southern side, consistent with the extinction being a possible cause behind the lack of a spiral-arm feature on the southern side. Additionally, we note that the optical colour is also redder for the southern side.  UV images thus clearly demonstrate a young star-forming disk in Speca.  From the UV color for the whole area of Speca (FUV-NUV $\\sim$ 0.43 in AB) and comparing with models of stellar population synthesis (Bruzual \\& Charlot 2003), we inferred that the star formation in Speca is younger than $\\sim$500 Myr. However, the $u'$-$r'$ colour (2.6 in AB), five band SED and nuclear spectrum from SDSS suggest a several Gyr old stellar population dominant in the host galaxy.  The surface brightness profile of the Lulin $R$-band image was analyzed with GALFIT version 2.03b (Peng et al. 2002). We explored both the exponential disk and de Vaucouleurs profile, with and without a nuclear Gaussian source. The best model was found to be an exponential disk model (effective radius = 5.3 arcsec, axial ratio = 0.35, P.A.=17$^{\\circ}$) with an unresolved (within the errors) point source at the nucleus. The radial profile of the surface brightness is shown in Figure 1 (d) with our best model and the next best one (de Vacouleurs profile without a nuclear source). The observed profile shows excess features on both the NE and SW sides of the nucleus at $\\sim$7 arcsec from the nucleus, making the fit unsatisfactory in terms of statistics even with our best model. However, the best model of disk with a nuclear point source is far better than the second best one since the latter predicts too much flux in both the inner and outer parts of the radial profile.  In summary, the optical and UV structure and optical radial brightness profile strongly support the identification of Speca as a rare radio galaxy hosted in a young star forming spiral/disk galaxy.  To the best of our knowledge, this is the second spiral-host large radio galaxy, after the only well-known case of 0313-192, discovered by Ledlow, Owen \\& Keel (1998).  In light of recent H{\\sc i} studies finding large massive gas disks (Oosterloo et al. 2007) and UV studies finding young stellar disks in a large-fraction of early-type galaxies (Yi et al. 2005, Salim \\& Rich 2010), we speculate that the disk nature of Speca is due to recent star formation, possibly following a merger more than 500 Myr ago (e.g. NGC~3801; Hota et al. 2009, 2011). \\subsection{Episodic radio jet activity} \\label{radio} The possible episodic nature of this radio galaxy was first identified from comparisons of NVSS and FIRST images, where two prominent emission blobs on large scales seen in NVSS were found to be missing in FIRST. To investigate further we observed it with the GMRT at 325 MHz. We found typical `edge-brightened' radio lobe structure as in FR II type radio galaxies (Figure 2 top panel).  A similar radio structure is also seen in the higher frequency FIRST image, where the lobes have equal flux densities of 24 mJy (23.9 mJy for north-western (NW) and 23.8 mJy for south-eastern (SE)) within the errors. The projected linear size of the lobes, 307 kpc (peak to peak measured to be 128 arcsec aligned at P.A. of 306$^\\circ$), is typical of classical large radio galaxies. These components are hereafter referred to as `middle lobes' of Speca.  To our surprise, we also found two new peaks (2.24 mJy beam$^{-1}$ for NW peak and 1.34 mJy beam$^{-1}$ for SE peak) of radio emission with a peak to peak projected separation of $\\sim$50 kpc (22 arcsec) at a P.A. of 297$^{\\circ}$.  The total flux densities of these blobs are 4.52 and 2.56 mJy for the NW and SE respectively.  The position angle of the line joining these two peaks is $\\sim$10$^\\circ$ offset from that defined by the middle lobes. Higher resolution and higher frequency images are required to confirm these peaks as radio lobes, but in comparison with B0925+420 (Brocksopp et al. 2007), we believe them to be lobes and refer as `inner lobes'.  Interestingly, the inner lobes are located along the minor axis of the host galaxy but are larger than its optical extent (29 kpc) suggesting that the lobes have not formed due to interaction with the inter-stellar medium of the host galaxy (Figure 2 top panel).  A large diffuse emission blob to the south of the SE middle lobe, and almost connected, can be seen in Figure 2 (top panel).  It shows a centrally peaked structure and the outer boundary on its SE side is almost straight with a sharp drop in surface brightness.  The total extent of this diffuse emission is $\\sim$300 kpc, comparable to the full extent of the middle lobes. It shows an overall NE-SW elongation roughly orthogonal to the middle lobes. This emission blob was prominent in NVSS but was missing in FIRST.  Because of this diffuse nature and no apparent optical counterpart (see Section 3.3) hereafter we refer to it as the `SE outer relic lobe'.  Corresponding emission on the NW side was very faint, and so we present the smoothed 325 MHz image of a larger region of the sky in Figure 2 (bottom panel).  In addition to the SE outer relic lobe and middle lobes, an emission blob located NW of the middle lobes shows a north-south elongation similar to that seen in NVSS.  Detailed comparison of its structure from all the three images available (325 MHz with a beam of 9 arcsec (red contours), 325 MHz with a beam of 45 arcsec (thick black contours) and 1400 MHz with a beam of 45 arcsec (blue contours)) is shown in the inset image. Although the overall orientation matches, the location of the emission peaks does not match, demonstrating the diffuse nature of the blob, and hence explaining why it was missed in the high angular resolution (beam $\\sim$5 arcsec) FIRST data. This diffuse nature along with the lack of any potential optical counterpart (see Section 3.3) suggests that it is not an independent radio source.  Hereafter we refer to this as `NW outer relic lobe'.  These relic lobes are quite asymmetric in terms of their location with respect to the host galaxy, radio flux densities and orientation with respect to the inner and middle lobes.  There is no signature of any more diffuse emission apart from the lobes described here.  Note that a compact radio source seen on the SE corner of the image is an unrelated source with a clear optical counterpart.  The structure of these two relic lobes neither bear similarity with typical radio lobes nor with the partial shell-like structures of double radio relics found in clusters (e.g. Bagchi et al. 2006).  Hence we believe that both these diffuse radio blobs are due to an earlier episode of jet-activity of the AGN hosted in Speca.  However, we can not completely rule out an origin for this diffuse emission by some independent radio-loud AGN or other cluster related phenomena.  The largest angular separation between these two relic lobes (570 arcsec) corresponds to $\\sim$1.3 Mpc, classifying Speca as a giant (old) radio galaxy (see Schoenmakers et al. 2001).  These three distinct pairs of lobes are probably produced by three different episodes of radio-jet activity of the AGN in Speca.  If confirmed, this will be the second triple-double radio galaxy after B0925+420 (Brocksopp et al. 2007). Multiple episodes of jet activity may be correlated with the recent ($\\lapp$ 500 Myr) star formation seen in the host galaxy (see Section 3.1).  Since these relic lobes could be several 100s of Myr old and the few detections of them so far are believed to be the ``tip of the iceberg'' (Dwarakanath \\& Kale 2010), upcoming sensitive low frequency surveys (e.g. TGSS: TIFR GMRT Sky Survey at 150 MHz or surveys from LOFAR: Low Frequency Radio Array) will possibly discover many more examples to help understand the duty cycle of AGNs, black hole coalescence (Liu 2004), and the role of AGN-feedback in galaxy and cluster evolution (Croton et al. 2006). \\subsection{Revived relic lobes in a cluster} \\label{cluster} Speca belongs to a cluster of galaxies MaxBCG J212.45357-03.04237, having 14 members within $r_{\\rm 200}$ (Koester et al. 2007).  While spectroscopic redshifts are available for only five other members in NED (Figure 2 bottom panel), we searched the SDSS database for further possible member galaxies with 0.135 $<$ $z_{\\rm photo}$ $<$ 0.141 and within a radius of 8 arcmin ($\\sim$1 Mpc for the distance of Speca). We found nearly 60 members (marked by triangles in Figure 2 bottom panel), further supporting the presence of a cluster of galaxies. The maximum velocity difference of $\\sim$1000 km s$^{-1}$ between confirmed members also supports the presence of a cluster of galaxies around Speca.  Although X-ray emission from this region has not been detected by the ROSAT All Sky Survey, we calculate that a 3$\\sigma$ upper limit corresponds to a luminosity of $7 \\times 10^{43}$ erg s$^{-1}$ and it is still consistent with the presence of a hot intra-cluster medium around Speca.  There is no member galaxy located in the middle of the diffuse emission.  Furthermore, by searching in the SDSS and NED database without any redshift constraints, we do not see any extended optical source ($>$2 arcsec) brighter than -18 in absolute $r'$-band magnitude (assuming SDSS $z_{\\rm photo}$) and potential radio-loud AGN host galaxy (red in the colour magnitude diagram) within the region of the diffuse lobes.  To understand the physical properties we investigated the spectral index and polarisation of the relic lobes.  The two point spectral index of the NW relic, which has $S_{325}$ = 46 mJy (from GMRT) and $S_{1400}$ = 28.45 mJy (from NVSS), is very flat ($\\alpha$=-0.32, where $S_{\\nu} \\propto \\nu^{\\alpha}$).  There is no chance of missing flux in our measurements at 325 MHz with the GMRT, for this small lobe of nearly 3 arcmin size.  Although the spectrum of the relic lobe should be confirmed by UV-matched imaging at multiple frequencies, flattening is significant and unlikely have been caused by synthesis imaging issues or the 5-10 per cent absolute flux density calibration errors.  On the other hand if NVSS has missed flux, then the true spectral index will be even flatter than -0.32.  So, although the exact value of the spectral index ($\\alpha$=-0.32) may not be accurate, its true spectral index has to be, we believe, flatter than typical cluster radio relics or relic lobes, which are usually steeper than -1 (Ensslin et al. 1998, Dwarakanath \\& Kale 2009).  Therefore, we propose that such a flat spectral index of the NW relic lobe has been caused by particle re-acceleration process (Brunetti et al. 2008, van Weeren et al. 2010).  In NVSS, the SE middle lobe and the SE outer relic lobes are not well separated.  Their total flux density is 108 mJy at 1400 MHz. However, from the FIRST image we see only the SE middle lobe, and its flux density is 24 mJy.  Subtracting this value from the total flux density of the SE middle and the SE outer lobes we calculate the flux density for the SE outer relic to be 84 mJy at 1400 MHz.  From the high resolution GMRT measurements, the flux density of SE relic lobe at 325 MHz is 250 mJy.  The resulting two point spectral index of the SE relic lobe is -0.75, which is $\\sim$0.2 flatter than that of the middle lobes (-0.92 for NE and -0.98 for SW lobes between 1400 MHz (FIRST) and 325 MHz).  This spectral index of the SE relic is also flatter than what is expected for such diffuse relic lobes at a similar redshift (e.g. Dwarakanath \\& Kale 2009).  So, some degree of particle re-acceleration is also required in the SE relic lobe.  The SE relic lobe shows a sharp linear edge and drops in surface brightness on the edges.  The fraction of polarized flux density, measured in NVSS, is 19 per cent, strongly suggesting ordering of the magnetic field due to large-scale linear compression, like in shocks (Ensslin et al. 1998, van Weeren et al. 2010).  The orientation of these elongated outer relic lobes shows little relation with the jet direction of the middle lobes.  In the absence of an X-ray image, the structure of the relics can be compared with the galaxy distribution in the cluster.  All the five bright, spectroscopically confirmed members are seen to the NE of Speca and seem to be early-type galaxies based on their morphology and red optical colour ($u'$-$r'$).  Speca is the largest and most dominant member of the cluster but the distribution of all possible sixty member galaxies does not show a higher concentration around it (Figure 2 bottom panel).  This distribution suggests Speca to be in a dynamically young cluster, possibly related to either a merger or accretion from galaxy filaments.  A cluster-cluster merger seems to be unlikely a cause, as the distribution of member galaxies does not show a double peak.  Furthermore, the morphology of the relic lobes does not resemble a case where arc-like radio emission is seen perpendicular to the elongated X-ray emission or double-peaked galaxy distribution of the merging clusters (Bagchi et al, 2006, van Weeren et al. 2010).  Any cluster-cluster merger along the line of sight can also be ruled out since the relics are widely separated ($\\sim$1.3 Mpc) as seen in projection.  Therefore, we speculate that the relic lobes are most plausibly revived by shocks at the cluster outskirts due to cosmological accretion from galaxy filaments (Ensslin et al. 1998; Pfrommer et al. 2008). Numerical simulations predict two locations of such accretion shocks, one near the virial radius (R$_{vir}$) and the other at $\\sim$ 3R$_{vir}$ (Molnar et al. 2009 and references there in). Following the formula from Girardi et al. (1998), a rough estimate of the projected velocity dispersion ($\\sim$ 300 km s$^{-1}$) of the spectroscopic members leads to a R$_{vir}$ of $\\sim$ 850 kpc. Since re-accelerated relic lobes can be expected at diametrically opposite ends of the cluster, the expected separation is $\\sim$ 1.7 Mpc.  Given the unknown orientation of the relic lobes to the line of sight, the largest separation of 1.3 Mpc between the relic plasma is consistent with it tracing the accretion shock at the virial radius of this cluster.  Since the cluster is possibly dynamically young and Speca is the brightest cluster galaxy (BCG) with young star formation, as argued in similar BCGs (McNamara et al. 1996, Hicks, Mushotzky \\& Donahue 2010), we speculate that, alternatively to a galaxy merger origin, Speca could be accreting cold gas from the surrounding intra-cluster medium. Such young star-forming radio galaxies may be a `missing link', as they are rare in the nearby Universe but likely common at higher $z$ (Norris et al. 2007, Zirm et al. 2005).  Galaxy mergers and/or cold gas accretion are likely to be correlated with the multiple episodes of AGN jet activity we see in Speca.  Most UV-bright young star formation seen in early-type galaxies is interpreted as residual star formation or incomplete quenching of the star formation by past AGN feedback (Schawinski et al. 2007). If this is the case, then the multiple episodes of feedback seen in Speca with an unusual host galaxy are a unique laboratory to investigate the galaxy-black hole co-evolution process in action, and an opportunity for near-field cosmology studies."
    },
    "1107/1107.4268_arXiv.txt": {
        "abstract": "The second RAdial Velocity Experiment (RAVE) Data Release (DR2) derives $\\log g$ values but we present a simpler and cleaner method of identifying dwarfs and giants using only magnitudes, which does not require spectroscopic analysis.  We confirm the \\citet{Bilir06a} procedure  which estimates the number of dwarfs and giants via their positions in the $J-V$ two magnitude diagram by applying it to RAVE DR2. It is effective in estimating the number of dwarfs and giants at $J-H>0.4$ compared to RAVE's $\\log g$ values. For $J-H\\leq0.4$, where dwarfs and subgiants show a continuous transition in the $J$ magnitude histogram, we used the Besan\\c con Galaxy model predictions to statistically isolate giants. The percentages of giants for red stars and for the whole sample are 85\\% and 34\\%, respectively. If we add the subgiants, the percentage of evolved stars for the whole sample raises to 59\\%. For the first time in the literature, we analysed the effect of CHISQ on RAVE's $\\log g$ values (CHISQ is the penalised $\\chi^2$ from RAVE's technique of finding an optimal match between the observed spectrum and synthetic spectra to derive stellar parameters). Neither the CHISQ values nor the signal-to-noise ratio bias RAVE $\\log g$ values.  Therefore the method of identifying dwarfs and giants via the two magnitude diagram has been verified against an unbiased dataset. ",
        "introduction": "Large surveys such as the Two Micron All Sky Survey \\citep[2MASS,][]{2mass06}, Sloan Digital Sky Survey/Sloan Extension for Galactic Understanding and Exploration \\citep[SDSS/SEGUE,][]{York00, Yanny09} and the RAdial Velocity Experiment \\citep[RAVE,][]{Steinmetz06} give the opportunity to a researcher to investigate different topics. However, one needs to select the appropriate data from these surveys and to classify them according to his/her aim. RAVE provides spectroscopic data observed from the southern hemisphere. Radial velocities are available in the first data release \\citep[DR1,][]{Steinmetz06} and spectroscopic analyses to provide information on values of stellar parameters (temperature, surface gravity and metallicity) are also included in the second  data release \\citep[DR2,][]{Zwitter08} additional to the radial velocities. One can find the stellar parameters in the literature supplemented by stellar position, proper motion and photometric measurements from the Deep Near Infrared Survey of the Southern Sky \\citep[DENIS,][]{Fouque00}, {\\em 2MASS} \\citep{2mass06} and Tycho-2 \\citep{Hog00} surveys.  Examples of scientific use of such a dataset are described in \\citet{Steinmetz03}. They include the identification and study of the current structure of the Galaxy and of remnants of its formation, recent accretion events, as well as discovery of individual peculiar objects and spectroscopic binary stars. For example, kinematic information derived from the RAVE dataset has been used to demonstrate the presence of a dark halo in the Galaxy \\citep{Smith07}. \\citet{Veltz08} identified kinematical discontinuities in the direction to the Galactic poles which separate a thin disc, thick disc and a hotter component. \\citet{Seabroke08} searched for infalling stellar streams onto the local Galactic disk and found that it is devoid of any vertically coherent streams containing hundreds of stars. \\citet{Coskunoglu11} estimated Local Standard of Rest (LSR) values from RAVE DR3.  \\citet{Klement11} classified the dwarfs and giants in RAVE and used this classification in stellar stream detection.  Classification of stars in RAVE as dwarf and giant populations is also the topic of this paper. However, the procedure we use here is different. We use the positions of stars in the $J-V$ two magnitude diagram to estimate the number of dwarfs and giants of the sample considered. According to this procedure, stars which lie above the line fixed by \\citet{Bilir06a} are dwarfs and those below the same line are giants. We show that this procedure provides accurate estimations compared to RAVE DR2 $\\log g$ values. Also, it can be applied without the following considerations of \\citet{Klement11} and hence provides simplicity: {\\em i}) the gravity of a given luminosity class changes with colour, {\\em ii}) reddening of stars causes to a less clear separation of dwarfs from other luminosity classes in $\\log g$ distribution, {\\em iii}) the dwarf-giant separation gets better with increasing colour, {\\em iv}) the expected small fraction of dwarfs in combination with their higher $\\log g$ uncertainties might hinder a clear separation and {\\em v}) the $\\log g$ distribution has been shown to vary with Galactic latitude.  Our work is the first one in the literature to analyse the effect of CHISQ in the RAVE data. CHISQ is the penalised $\\chi^2$ from RAVE's technique of finding an optimal match between the observed spectrum and synthetic spectra to derive stellar parameters. It is conceivable that sample stars with large CHISQ and/or small signal-to-noise ratios (S/N) contaminate the dwarf and giant separation with unreliable $\\log g$ values. We investigate whether CHISQ or S/N biases RAVE $\\log g$ values and therefore biases the verification of our method of identifying dwarfs and giants via the two magnitude diagram.  The paper is organized as follows: the data and field dwarf-giant separation are given in Sections 2 and 3, respectively. Confirmation of field dwarf-giant separation and the effects of S/N and $\\chi^2$ on the separation of dwarfs and giants are presented in Sections 4 and 5 and finally a summary and discussion is devoted to Section 6. ",
        "conclusions": "We estimated the number of stars for different populations in the RAVE DR2 by means of their positions in the $J$ versus $V$ two magnitude diagram \\citep{Bilir06a} and the Besan\\c con Galaxy model \\citep{Robin03}. The procedure of \\citet{Bilir06a} is scaled for separation of dwarfs and giants (but not for subgiants). One can predict the number of subgiants, additional to dwarfs and giants, via the Besan\\c con Galaxy model. Different results were obtained for two samples of stars, i.e. for blue and red ones, separated by $J-H=0.4$.  {\\bf a) Sample of stars with $J-H>0.4$:} Dwarfs and giants in this sample have been separated in a very simple manner, i.e. by their positions in the $J-V$ two magnitude diagram. No additional constraints are needed for this separation. The percentages of dwarfs (15\\%) and giants (85\\%) are confirmed by the percentages of stars predicted by the Besan\\c con Galaxy model, i.e. 18\\% dwarfs, 82\\% giants and no subgiants. A second confirmation was done using the work of \\citet{Ammons06}. The positions of dwarf and giant in \\citet{Ammons06} occupy the dwarf and giant half planes in our work. Finally, the positions of the stars with $J-H>0.4$ in the $\\log g$ versus $J-H$ diagram indicate a large number of giants. Actually, most of these red stars have surface gravities less than the upper limit attributed to giants, i.e. $\\log g=3.8$ \\citep{Zwitter10}.  {\\bf b) Sample of stars with $J-H\\leq0.4$:} The separation of stars with $J-H\\leq0.4$ by means of their positions in the $J$ versus $V$ two magnitude diagram is not as easy as for stars with $J-H>0.4$, because there is a considerable number of subgiants among these stars which were not scaled in \\citet{Bilir06a}. The procedure of \\citet{Bilir06a} reveals 82 giants and 6877 dwarfs corresponding to 1\\% and 99\\% of the total number of stars, respectively. However, the percentage number of giants, dwarfs and subgiants predicted by the Besan\\c con model \\citep{Robin03} are 2\\%, 65\\% and 33\\%, respectively. Hence, we plotted the histogram for $J$ magnitude, for stars with $J-H\\leq0.4$ to investigate this problem. The unique modality of the distribution indicates that there is a continuous transition between dwarfs and subgiants. Hence, we applied the Besan\\c con Galaxy model to this sample for estimation of the number stars for different populations. Thus, the number of dwarfs reduced to 58\\%, and the remaining 42\\% of the total sample turned out to be evolved stars (giants and subgiants).  It should be noted that dwarfs and giants with $J-H\\leq0.4$ in \\citet{Ammons06} confirm the dwarf and giant half planes, respectively, defined by $J$ and $V$ magnitudes. The identification of dwarfs and giants by \\citet{Ammons06} is based on atmospheric parameters. Hence, its accuracy is high.  {\\bf c) Error in $\\log g$:} 18\\% of the total stars with $3.5<\\log g<4$ in Fig. 5 occupy the subgiant segments of the isochrones. Hence, they should be subgiants. However, the  percentages of subgiants predicted by the Besan\\c con model is 25\\%. The difference between percentages can be explained by the estimated error on RAVE DR2 surface gravities being as large as 0.5 dex.  {\\bf d) Effect of S/N and CHISQ values:} We showed in this work that neither the small signal to noise ratio (S/N) nor the CHISQ values bias RAVE $\\log g$ values. Therefore the method of identifying dwarfs and giants via the two magnitude diagram has been verified against an unbiased dataset.  {\\bf e) Comparison between the results in our work and the ones apparent in the literature:} The first estimate of the percentage of different stellar populations in RAVE was published in \\citet{Seabroke08}. They used a reduced proper motion diagram to kinematically separate red ($J-K>0.5$) giants from red ($J-K>0.5$) dwarfs. Statistically, they found that 44\\% of their RAVE sample consisted of red ($J-K>0.5$) giants and red ($J-K>0.5$) subgiants. Although \\citet{Seabroke08} do not explicitly state their 44\\% sample includes subgiants, their fig. 13 shows the subgiant branch extends to redder colours than $J-K=0.5$. \\citet{Seabroke08} do not distinguish between blue ($J-K<0.5$) subgiants and blue ($J-K<0.5$) dwarfs. Therefore \\citet{Seabroke08} do not provide an estimate for the percentage of evolved (subgiant and giant) stars in RAVE. Given their 44\\% sample excludes blue ($J-K<0.5$) subgiants, 44\\% represents a lower limit on the percentage of evolved stars in RAVE, which is consist with the value of 59\\% that we find in our work. It should also be noted that the \\citet{Seabroke08} RAVE sample consists of many more stars than DR2 ($\\sim$200~000 stars compared to $\\sim$50~000) so the density of sampling in different directions on the sky will be different between the two samples and so they are not suitable for detailed direct comparison. \\citet{Klement11} found 43\\% of RAVE DR2 are dwarfs (see completeness column in their table 3). The remaining 57\\% are subgiants, giants and supergiants, which agrees closely with our value.  \\begin{table} \\setlength{\\tabcolsep}{4pt} \\centering \\caption{Comparison of percentages of dwarfs and giants with different ranges of S/N.} \\begin{tabular}{ccccccccccc} \\hline & \\multicolumn{5}{c}{Number of dwarfs} & \\multicolumn{5}{c}{Number of giants}\\\\ \\hline Subsample & Total&  S/N$\\leq$33 & \\% & S/N$>$33 & \\% & Total&  S/N$\\leq$33 & \\% & S/N$>$33 & \\% \\\\ \\hline $\\log g\\leq 4$, $J-H\\leq0.4$ & 2520 & 444 &  18 & 2076 & 82 & 31   & 2   & 7  & 29   & 93\\\\ $\\log g\\leq 4$, $J-H>0.4$    & 484  & 112 &  23 & 372  & 77 & 3567 & 316 & 9  & 3251 & 91\\\\ $\\log g>4$, $J-H\\leq0.4$     & 4357 & 939 &  22 & 3418 & 78 & 51   & 11  & 22 & 40   & 78\\\\ $\\log g>4$, $J-H>0.4$        & 211  & 51  &  24 & 160  & 76 & 249  & 39  & 16 & 210  & 84\\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\setlength{\\tabcolsep}{4pt} \\centering \\caption{Comparison of percentages of dwarfs and giants with different ranges of CHISQ ($\\chi^2$).} \\begin{tabular}{ccccccccccc} \\hline & \\multicolumn{5}{c}{Number of dwarfs} & \\multicolumn{5}{c}{Number of giants}\\\\ \\hline Subsample & Total&  $\\chi^2$$\\leq$2290& \\% & $\\chi^2$$>$2290& \\% & Total&  $\\chi^2$$\\leq$2290& \\% & $\\chi^2$$>$2290& \\% \\\\ \\hline $\\log g\\leq 4$, $J-H\\leq0.4$ & 2520 & 1875 & 74 & 645 & 26 &   31 &   16 & 52 &   15& 48\\\\ $\\log g\\leq 4$, $J-H>0.4$    &  484 &  414 & 86 &  70 & 14 & 3567 & 2083 & 58 & 1484& 42\\\\ $\\log g>4$, $J-H\\leq0.4$     & 4357 & 3541 & 81 & 815 & 19 &   51 &   40 & 78 &   11& 22\\\\ $\\log g>4$, $J-H>0.4$        &  211 &  177 & 84 &  34 & 16 &  249 &  178 & 71 &   71& 39\\\\ \\hline \\end{tabular} \\end{table}  As conclusion, stars with $J-H>0.4$ can be separated into dwarf and giant populations in a very simple manner, i.e. by their positions in the $J-H$ two magnitude diagram. One does not need any other constraints for such a separation. Our argument has been confirmed by the work of \\citet{Ammons06} and the Besan\\c con Galaxy model \\citep{Robin03}. However, the existence of a considerable number of subgiants complicates the separation of stars with $J-H\\leq0.4$ into different population types. The difference between the percentage of subgiants obtained via the $\\log g$ versus $J-H$ diagram (18\\%) and the one estimated by using the Besan\\c con Galaxy model (25\\%) supports the values in $\\log g$ of RAVE dataset. The percentage of evolved (subgiant and giant) stars found in this work is consistent with \\citet{Seabroke08} and in close agreement with \\citet{Klement11}. For the first time in the literature, we analysed the effect of the CHISQ in the RAVE data.  Neither the CHISQ values nor the signal-to-noise ratio bias RAVE $\\log g$ values.  Therefore the method of identifying dwarfs and giants via the two magnitude diagram has been verified against an unbiased dataset."
    },
    "1107/1107.3024_arXiv.txt": {
        "abstract": "By using recent data, we directly determine the dark matter (DM) induced $e^\\pm$ spectrum at the source from experimental measurements at the earth, without reference to specific particle physics models. The DM induced gamma rays emitted via inverse Compton scattering are then obtained in a model independent way. However the results depend on the choice of the astrophysical $e^\\pm$ background, which is not reliably known. Nevertheless, we calculate, as an illustration, the fluxes of gamma rays from the Fornax cluster in the decaying DM scenario with various astrophysical $e^\\pm$ backgrounds. Without any assumptions on details of the DM model, the predictions turn out to be either in disagreement with or only marginally below the upper limits measured recently by the Fermi-LAT Collaboration. In addition, these DM induced ICS gamma rays in the GeV range are shown to be almost independent of choices of cosmic ray propagation model and of DM density profile, when a given astrophysical $e^\\pm$ background is assumed. This provides a strong constraint on decaying DM scenario as the gamma rays may be produced in other processes besides inverse Compton scattering, such as the bremsstrahlung and neutral pion decays. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.4953_arXiv.txt": {
        "abstract": "I show that, in the context of MOND, non-isothermal models, approximated by high order polytropic spheres, are consistent with the observations of the radial distribution of the line-of-sight velocity dispersion in the distant globular cluster, NGC 2419.  This calls into question the claim by Ibata et al. that the object constitutes a severe challenge for MOND. In general, the existence and properties of globular clusters are more problematic for LCDM than for MOND. ",
        "introduction": "In a recent preprint, Ibata et al. (2011) have claimed that in the distant globular cluster, NGC 2419, the observed radial profile of line-of-sight stellar velocity dispersion, combined with the surface density distribution of starlight, is consistent with Newtonian gravity in the context of a broad class of globular cluster models and inconsistent with modified Newtonian dynamics (MOND). They go further and assert that this object is a ``crucible'' for acceleration-based theories of gravity primarily because its large galacto-centric distance means that the galactic external field is less than 20\\% of the MOND critical acceleration, $a_0$.  Thus, the complications due to the MONDian external field effect should be negligible; it is a relatively isolated system and can be modeled taking only internal dynamics into account.  I argue here that the elevation of this object to the role of crucible is overstated primarily due to the limitations of the models considered and that a class of MONDian models which deviate slightly from an isothermal state are consistent with the observations.  More generally, I claim that the very existence and overall properties of globular clusters are more problematic for LCDM than for MOND. ",
        "conclusions": ""
    },
    "1107/1107.2602_arXiv.txt": {
        "abstract": "A mechanism capable to provide a natural solution to two major cosmological problems, i.e. the cosmic acceleration and the coincidence problem, is proposed. A specific brane-bulk energy exchange mechanism produces a total dark pressure, arising when adding all normal to the brane negative pressures in the interior of galactic core black holes. This astrophysically produced negative dark pressure explains cosmic acceleration and why the dark energy today is of the same order to the matter density for a wide range of the involved parameters. An exciting result of the analysis is that the recent rise of the galactic core black hole mass density causes the recent passage from cosmic deceleration to acceleration. Finally, it is worth mentioning that this work corrects a wide spread fallacy among brane cosmologists, i.e. that escaping gravitons result to positive dark pressure. ",
        "introduction": "During last decades it has been realized that the investigation of the problems associated with the cosmological constant would provide an insight into the structure and the properties of elusive quantum gravity. A serious problem concerning the cosmological constant refers to the vast discrepancy between the value a theorist would expect and the very low value of the effective cosmological constant. As Zel'dovich \\cite{zeld} first noticed, the effective cosmological constant we measure is the sum of the pure geometric origin cosmological constant plus the energy density of the vacuum. It seems impossible to understand why the measured effective cosmological constant is so much smaller than the value of the vacuum energy calculated by a quantum field theorist (cosmic phase transition, quantum field zero-point energies). This puzzle challenges the inflationary scenario and the various models of quantum gravity.  The present work attempts to solve a recently emerged problem regarding the cosmological constant issue which concerns the measured cosmic acceleration. During the last decade, it has been established through different independent pieces of astronomical data that empty space, devoid of the usual matter, is anti-gravitating. It creates gravitational repulsion and gives rise to an accelerated cosmological expansion. According to our present-day understanding, this accelerated expansion could be induced by an effective cosmological ``constant\"-like term (vacuum $p=-\\rho $ or dark energy $p<-\\frac{\\rho }{3}$). Data suggest that the magnitude of the required vacuum/dark energy is quite close to the critical (closure) cosmological energy density (approximately 70\\%). Why vacuum energy, which stays constant in the course of cosmological evolution, or why dark energy, which evolves with time quite differently from the normal matter, have similar magnitude with matter density just today, all being close to the value of the critical energy density?  There are several ideas in literature, though yet incomplete, that have the potential to provide solutions to cosmic acceleration problem. The present paper proposes that \\emph{the negative five-dimensional pressure produced from an astrophysical brane-bulk energy outflow occurring inside all cosmic black holes} is large enough to drive the measured cosmic acceleration. Assuming an RS-like cosmological brane \\cite{RS}, \\cite{Roy}, \\cite{gregory}, this total pressure (called dark pressure) arising from the sum of all negative pressures normal to the brane suffices to explain the coincidence problem, without using a negative vacuum energy or exotic fields/fluids throughout the universe. In the presented scenario dark energy ``originates\" from dark matter (grows from a negligible value to a significant one due to dark pressure); moreover, the recent appearance of the cosmic acceleration is correlated to the recent increase of the galactic core black hole density. An accelerating universe has already been produced by several brane-bulk energy exchange scenarios \\cite{zarikas}, \\cite{tamva}, \\cite{tetradis}, \\cite{langlois}, \\cite{Hebecker}. However, in all these works the whole universe should be hot enough, and therefore, these scenarios fail to explain recent acceleration.  Both astrophysical black holes in haloes and supermassive black holes at the galactic centres appear after the large scale structure of the universe, weight a portion of $\\rho _{m}$ and are regions where high energy interactions occur. This fact will be at the center of the proposed mechanism. Assuming that a brane cosmological model describes our universe, it is natural to expect a moderate exchange of energy between the brane and the bulk. Astrophysical black holes contain matter in an unknown form, i.e. effective quantum fluid (possibly arising from superposition of non empty black hole quantum spacetimes) and accrete continuously mass. Collapsing matter falling into a black hole accelerates, interacts and gets easily ``thermalized\" to temperatures close and above $M$ (five-dimensional fundamental Planck mass). Furthermore, it is expected portion of black hole mass to be in the form of highly energetic states close to $M$, not only due to accreting matter interactions but also due to Hawking-like particle production in the interior \\cite{greenwood}. But for energy scales close to $M$, matter interactions result to graviton escape to the bulk. Therefore, energy outflow can occur in the interior of galactic halo black holes and galactic core supermassive black holes. This black hole originated exchange results to non-zero energy-momentum tensor components $T_{05}$ and $T_{55}$ and is able to provide the necessary conditions for a cosmic acceleration.  The emission of gravitons to the bulk is associated with negative dark pressure on the brane. Indeed, a brane experiences positive pressure when bulk particles fall into it. This negative pressure is the responsible quantity that provides the required amount of the measured cosmic acceleration. Although there is a small outflow in our scenario in each of the galactic black holes (consistent with black hole mass evolution and galaxy dynamics), this leakage is associated unavoidably with orders of magnitude larger dark pressure. Moreover, the emerged dark energy is not a small portion of the dark matter but of the same order with it. Note that in our scenario, before outflow starts (before galaxy formation) there may be either a zero or a very small positive non zero decelerating dark radiation term $\\frac{\\mathcal{C}}{a^4}>0$. However, this radiation term overpasses well known problems of nucleosynthesis constraints \\cite{ichicki}. When negative dark pressure emerges, this radiation term is modified and finally becomes an accelerating dark energy component.    Note also that it costs almost nothing to stretch a brane that has zero total tension/cosmological constant \\cite{thooft}. Apparently spacetime is such that it takes a lot of energy to curve it, while stretching it is almost for free, since the cosmological constant is zero. This is quite contrary properties of objects from every day experience, where bending requires much less energy than stretching. The present study tries only to explain naturally the recent cosmic acceleration while assumes that there is some mechanism that sets the cosmological constant from field theory vacuum energy equal to zero (for example R-S fine tuning, holography etc.).  The paper is organized as follows. In section II the mathematical framework is presented. Here, it becomes obvious from Eq. (\\ref{simple2}) that the negative dark pressure $\\Pi$ can drive acceleration. In addition, Eq. (\\ref{systemb3}) shows that $\\Pi$ determines also the time derivative of dark energy, resulting to the current value of $w_{DE}$. Furthermore, the question of how large should be the present value of dark pressure $\\Pi$, is estimated in Eq. (\\ref{current po}). Such a value can easily be obtained, as it is explained in sections III and V, for moderate values of the involved astrophysical parameters. In section III the connection of the mathematical framework with the astrophysical context is given. Section IV provides various supportive theoretical aspects of the physics of the proposed mechanism. However, section IV is not crucial to the main results of the analysis.  In section V numerical computations are presented in order to prove the success of the model. Results have been derived both analytically and numerically (for verification reasons) and are presented together with descriptions of the range of the involved parameters that ensure: i) recent passage from deceleration to acceleration, ii) small outflow that do not violate black hole evolution/mass, iii) nucleosynthesis bounds, and iv) universe age. Finally, section VI is dedicated to the conclusions. ",
        "conclusions": "This work proposes that the recent cosmic acceleration happens as a result of the negative five-dimensional pressures produced in the galactic black holes of a brane cosmological RS setup. It was proved that the total contribution $\\Pi$ from all the dark pressures from all the centers of galaxies suffice to provide the measured cosmic acceleration. An exciting outcome is that the recent passage from the deceleration to the acceleration era happens due to the recent increase of the galactic core black hole mass density. Based on the derived expressions we have shown that it is easy to get the expected negative values of cosmic acceleration $w_{DE,0}\\simeq -1$ and dark energy $\\Omega_{DE,0}=0.7$ \\emph{even for conservative values of all the relevant parameters}, i.e. for small values of the mean temperature $\\Theta _{mean}$ in the interior of the black holes, for small values of galactic core black holes masses and for values of the five-dimensional Planck mass $M$ several decades of TeVs. In our mechanism, outflow is associated unavoidably with a large dark pressure. The magnitude of the produced dark pressure is connected with that of dark radiation through the equation of state of the quantum fluid in the interior of a black hole. Qualitatively one can immediately check the efficiency of producing the observed cosmic acceleration estimating the required amount of dark pressure shown in equations (\\ref{q0}), (\\ref{current po}) or (\\ref{current po2}). This value of dark pressure can naturally be realized in our case based on the known astrophysical data and the assumed temperature of the effective fluid. Of course, in order cosmic acceleration to be proven, the system of differential equations (\\ref{systemc1}), (\\ref{systemc2}) has to be solved, and indeed it is seen that the energy exchange along with the dark pressure give an order one effect on cosmological scales.  The proposed mechanism has several advantages: i) it is independent of the bulk matter and consequently retains predictability, ii) the associated values of $\\psi$ in the Hubble evolution (\\ref{systemb2}) originate from the brane black hole astrophysical phenomenon of energy outflow $T$ and its associated pressure $\\Pi $ along the fifth dimension, and not from the motion or the position of the brane in the bulk, thus again retaining predictability, iii) the mechanism is \\textquotedblleft on\" at present times and \\textquotedblleft off\" at the early stages of the cosmic evolution explaining naturally coincidence problem, iv) it relates the amount of the produced acceleration with the present matter content, and v) it produces easily cosmic acceleration for sensible values of the relevant parameters.  However, the most interesting and worth mentioning finding is the fact that sensible and safe values of outflow, as those appeared in Table 2, suffice to result to the observed cosmic acceleration. A first reason behind this outcome is the large number of galaxies in the universe. Another reason is that the additive effect of all outflows and associated much larger dark pressures from each galactic center result to a non negligible kinetic effect, i.e. acceleration due to the geometry of the setup. Finally, the dark pressure drives towards acceleration from earlier times of the cosmic evolution and not just today. Since energy density of the galactic core black holes increases as redshift decreases at recent times ($z<2$), dark pressure becomes stronger driving the passage to the acceleration era.  In summary, the novelties of the present article are: 1) correction of a wide spread mistake among the brane cosmologists that outflow is associated with a positive and not negative dark pressure, 2) the presentation of a new astrophysical origin mechanism of brane-bulk energy exchange, 3) new braneworld solutions describing the evolution of brane equations with non zero $T_{05}$, $ T_{55}$, 4) new gravitational collapse solution on a brane with non zero $T_{05}$, $T_{55}$, 5) numerical results estimating the produced cosmic acceleration, and 6) correlation of the measured acceleration to the recent rise of the galactic core black hole cosmic energy density.  The calculations of the scenario could have failed for various reasons: if a very high temperature was needed in the interior of the black hole, or if a small fundamental Planck scale or a large $T^{0}_{5}$ was needed conflicting of course with galactic dynamics, or if for the given variation of the cosmic black hole density as a function of redshift the cosmic evolution failed to posses a long deceleration era accompanied by a recent acceleration one. However, the concrete and conservative numerical values used lead the scenario to success.    One interesting point worth to be raised is that the proposed mechanism could work together with various studies that suggest solutions to the cosmological problem. For example, there are recent holographic ideas capable to explain why the cosmological constant should be almost zero \\cite{thomas}, \\cite{hsu}. However, it appears to have a difficulty to explain naturally the cosmic equation of state. Therefore, the present work together with all type of holographic explanations can provide a complete solution to the general problem of the cosmological constant value."
    },
    "1107/1107.5214_arXiv.txt": {
        "abstract": "We analyse the CMB radiation in spherical 3-spaces with non-trivial topology. The focus is put on an inhomogeneous space which possesses observer dependent CMB properties. The suppression of the CMB anisotropies on large angular scales is analysed with respect to the position of the CMB observer. The equivalence of a lens space with a Platonic cubic space is shown and used for the harmonic analysis. We give  the transformation of the CMB multipole radiation amplitude as a function of the position of the observer. General sum rules are obtained in terms of the squares of the expansion coefficients for invariant polynomials on the 3-sphere. ",
        "introduction": "\\label{sec:intro}  Cosmic topology examines multi-connected manifolds as candidates for the spatial part of cosmic space-time. It simulates the cosmic microwave background (CMB) radiation using the eigenmodes of the multi-connected manifolds, and explores if the specific multipole amplitudes and selection rules arising from the manifold are encoded in the CMB radiation.  In this paper the focus is put on spherical spaces. The simply connected 3-sphere underlies Einstein's initial cosmological analysis of 1917 \\cite{Einstein_1917}. In that year de Sitter \\cite{deSitter_1917} already discussed the projective space $\\mathbb{P}^3$ as an alternative to Einsteins model. This was the first cosmological application of a spherical model which is multiply connected, i.\\,e.\\ being a closed three-dimensional piece of Einstein's 3-sphere.  The topology of a manifold ${\\cal M}$ is locally described by homotopy. Homotopy composes loops on the manifold by concatenation and explores the homotopy group $\\pi_1({\\cal M})$ formed by these loops. If any loop can be continuously contracted to a point, the manifold is simply connected. Any topological manifold has an image on a simply connected covering manifold. The covering manifold $\\tilde{{\\cal M}}$ is tiled by copies of ${\\cal M}$. Covering manifolds and their tiles can be spherical, Euclidean or hyperbolic. In terms of the Riemannian metric they display constant positive, zero or negative curvature. In this paper, we concentrate on spherical manifolds which tile the 3-sphere ${\\mathbb S}^3$.  The second concept of a cover offers a global and quasi-crystallographic view of a topological manifold. On the covering manifold $\\tilde{{\\cal M}}$ there is a group ${\\rm deck}({\\cal M})$ of deck transformations which by fix-point free action tiles the covering manifold. Seifert and Threlfall \\cite{SE34} show that the two groups are isomorphic, $H={\\rm deck}({\\cal M})\\sim\\pi_1({\\cal M})$, and so the first local and the second global view of topology, in terms of the group $H$ as topological invariant, are equivalent.  In an abstract approach, a topological manifold is taken as a quotient space $\\tilde{{\\cal M}}/H$ of the covering manifold $\\tilde{{\\cal M}}$ by the group $H$. If this notion refers to the abstract group and not to a representation thereof, it leaves open the geometric form of the manifold. An algebraic characterisation for  the classification of spherical space forms is given by Wolf \\cite{WO84} in terms of unitary matrix representations of groups $H$, acting on the 3-sphere.   An introduction into the topological concept applied to the cosmological framework can be found in \\cite{LachiezeRey_Luminet_1995,Levin_2002} where all three spatial curvatures are discussed. The focus was shifted to spherical spaces \\cite{Gausmann_et_al_2001} by the paper \\cite{LU03} which claims that the low power in CMB anisotropies at large scales can be described by the Poincar\\'e dodecahedral topology. Thereafter, a lot of papers discussed the relevance of this result, and other spherical spaces such as the truncated cube and the tetrahedral space were investigated with respect to their statistical CMB properties, see e.\\,g.\\ \\cite{Roukema_et_al_2004,Gundermann_2005,AU05,AU05b,AU06,% Lustig_2007,Roukema_Kazimierczak_2011}. In addition, spherical lens spaces $L(p,q)$ are studied in \\cite{Uzan_et_al_2003}. The fundamental domain can be visualised by a lens-shaped solid where the two lens surfaces are identified by a $2\\pi q/p$ rotation for relatively prime integers $p$ and $q$ with $0<q<p$. For more restrictions on $p$ and $q$, see  below and \\cite{Gausmann_et_al_2001}. A further family of spherical spaces, the so-called Platonic polyhedra, were constructed from their homotopy groups and studied in \\cite{KR05,KR10,KR10B,KR10C}.  In the present paper, we examine the equivalence of spherical manifolds, the observer dependence of the multipole expansion, and quadratic sum rules following from the reduction of representations. We show the equivalence of the Platonic cubic manifold $N2$ and the lens manifold $L(8,3)$ and investigate the statistical properties of CMB anisotropies. The geometry of the manifold can be expressed by the Voronoi domain (see below) which has the observer of the CMB radiation in its centre. It turns out that the Platonic cubic geometry of the manifold $N2$ is equivalent to a special observer position in the manifold $L(8,3)$. This manifold is not homogeneous and is thus called inhomogeneous. In order to emphasise this point, consider two observers where the first observer position can be mapped by a transformation $M$ onto the second one. Assume that the first observer determines his Voronoi domain by the group elements $g\\in H$, then the second observer gets his Voronoi domain by the group $M g M^{-1}$. That is a similarity transformation, or a coordinate transformation. If $M$ and $g$ commute for all $g\\in H$, then both observers see the same Voronoi domain. In this case, one has a homogeneous manifold. On the other hand, if $M$ and $g$ do not commute, one obtains observer dependent Voronoi domains and thus an inhomogeneous manifold. The interesting point of view with respect to cosmic topology is that the statistical properties of the CMB radiation is observer dependent in the inhomogeneous case. As two examples for homogeneous manifolds, we also analyse the lens space $L(8,1)$ and the Platonic cubic manifold $N3$. The latter is equivalent to a manifold generated by the binary dihedral group $D_8^*$ isomorphic to the quaternion group $Q$. ",
        "conclusions": "In this paper we study spherical models of our Universe which possess spatial spaces that are multi-connected. These spaces with positive spatial curvature arise by tiling the 3-sphere ${\\mathbb S}^3$ using a deck group $H$. We restrict attention to groups of order 8 such that the volume of the considered manifolds is $\\hbox{vol}({\\mathbb S}^3)/8 = 2\\pi^2/8$. There are two such lens spaces $L(8,1)$ and $L(8,3)$, and the $D_8^*$ manifold which is obtained by the binary dihedral group $D_8^*$. Furthermore, there are two cubic Platonic manifolds $N2$ and $N3$, however, it is shown that $N2$ and $L(8,3)$ are equivalent and $N3$ corresponds to $D_8^*$.  Such a multi-connected space can be represented by its fundamental domain which is the subset of points such that no element of $g \\in H,\\: g\\neq e$ can operate  inside the domain, but any point of the cover outside the domain can be reached by the action of $H$ on a point inside the domain. Such a cell is called a Voronoi domain. It is the natural domain for an observer sitting at the origin of the coordinate system for which the group elements $g\\in H$ are defined. If the group elements $g\\in H$ are independent of the choice of the position of the observer, all observers construct the same Voronoi domain. Such manifolds are called homogeneous. On the other hand, if there is such a dependence, the shape of the Voronoi domain varies and such manifolds are called inhomogeneous. From the above manifolds, only $L(8,3) \\equiv N2$ is inhomogeneous. One can construct a spatial curve along which the observer can be shifted such that at one position the lens shaped Voronoi domain emerges whereas at an other position, the cubic Platonic Voronoi domain $N2$ appears. Other positions do not lead to domains having special properties.  The important point for cosmic topology is that the statistical properties of the CMB anisotropies depend on the group elements $g\\in H$ defined for an observer that sits at the origin of the coordinate system. Thus, inhomogeneous manifolds allow a much richer variety of CMB anisotropies, and the comparison with observations is much more involved. For the multi-connected spherical spaces generated by groups $H$ of order 8, the temperature 2-point correlation function $C(\\vartheta)$, eq.\\,(\\ref{Eq:C_theta}), and the multipole moments $C_l$, eq.\\,(\\ref{Eq:Cl_ensemble_general}), are computed for the CMB anisotropies. A suitable measure for the suppression of the anisotropies at large angular scales is the $S(60^{\\circ})$ statistics (\\ref{Eq:S_60}) which emphasises the unusual behaviour at scales $\\vartheta\\geq 60^{\\circ}$. The focus is put on the inhomogeneous space $L(8,3) \\equiv N2$, for which a one dimensional sequence of observer positions is derived which exhaust all possibilities allowed by this inhomogeneous space. The CMB anisotropies are calculated for this sequence, and it turns out that the strongest large scale suppression occurs for the Platonic cubic Voronoi domain, which has the observer in the centre of the $N2$ cell. On the other hand, the least suppression is observed for the lens shaped domain. This is in agreement with the hypothesis that well proportioned domains give the largest suppression of power. Thus, the example of the inhomogeneous lens space $L(8,3)\\equiv N2$ demonstrates that it can be premature to classify all lens spaces as not-well proportioned domains and thus uninteresting with respect to the observed CMB power suppression. Nevertheless, for $\\Omega_{\\hbox{\\scriptsize tot}}>1.07$, the homogeneous Platonic space $N3$ leads to an even stronger suppression of large scale CMB anisotropy than it is the case for all observer positions in $N2$.  The special case that appears when the distance $\\tau_{\\hbox{\\scriptsize SLS}}$ to the surface of last scattering satisfies $\\tau_{\\hbox{\\scriptsize SLS}} = \\pi/2$ is studied, although the cosmological parameters are unrealistic in this case. Then the topology requires an antipodal symmetry which partly survives in the CMB anisotropies and is reflected in the $S(60^{\\circ})$ statistics. It is discussed in detail how the different contributions affect the CMB signal in distinct ways as presented in fig.\\,\\ref{fig:spectrum_P3}."
    },
    "1107/1107.2434_arXiv.txt": {
        "abstract": "Motivated by observations that only a very few stars have been found to have antisolar differential rotation, much weaker in amplitude than that of the Sun, we analyze the stability of antisolar and solar type latitudinal differential rotations in the tachoclines of typical F,G and K stars. We employ two three-dimensional thin shell models, one for a Boussinesq but non-hydrostatic system, the other for a hydrostatic but non-Boussinesq system. We find that, in general, the combination of toroidal field band and differential rotation is more unstable, and unstable for lower toroidal fields, for antisolar than for solar-type differential rotation. In the antisolar case, the instability is always found to weaken the differential rotation, even if the primary energy source for the instability is the magnetic field. This favors surface antisolar differential rotations in stars being weaker than solar types, if the instability in the tachocline is felt at the surface of the star. This is most likely to happen in F stars, whose convection zones are much thinner than they are in G and K stars. This effect could help explain why the antisolar differential rotations that have been found are very weak compared to the Sun. ",
        "introduction": "The internal rotation pattern of a star has most extensively been studied observationally and theoretically for our nearest G-star, the Sun. For the surface of the Sun, \\citet{carrington1858} showed that the equator rotates faster than the pole, by systematically monitoring the rotation of sunspots over the 11-year cycle. In the 20th century, the Sun's positive pole-to-equator differential rotation was established from Doppler shift measurements \\citep{delury1939, plaskett1959}. A review of early observations of solar surface differential rotation can be found in \\citet{gilman1974}. Quantitative analyses of long term Mount Wilson Observatory data indicates that the equator rotates about 132 nHz faster than the pole \\citet{uetal1988} at the surface. Helioseismic measurements revealed that the Sun's positive pole-to-equator differential rotation pattern persists from the surface down to the core-envelope interface, through the bulk of the convective envelope -- the details can be found in an extensive review by \\citet{tcmt2003}.  Non-axisymmetric instability of solar-like, positive pole-to-equator differential rotation in the presence or absence of magnetic fields have been extensively studied in global 2D, quasi-3D and 3D thin-shell models \\citep{watson1981, dk1987, gf1997, dg1999, cdg1999, garaud2000, dg2001, cally2001, gd2002, cdg2003, dgr2003, cally2003, dcg2004, asr2005, asr2007, mgd2007, gdm2007}. Most of these calculations were applied to the solar tachocline, the interface between the solar convection zone and interior, where the stratification is thought to be either slightly subadiabatic, as in the overshoot layer where convection penetrates from above, or much more subadiabatic, as in the radiative layer immediately below the overshoot layer. But differential rotation anywhere in a star may be subject to the instabilities described in the above references.  Axisymmetric instability of latitudinal differential rotation can not occur in 2D HD or MHD (see discussions in \\citet{gf1997, dg1999, cdg2003}) and is hard to excite in quasi-3D shallow-water models for which about 2 million Gauss peak toroidal magnetic field is required. In a 3D thin-shell model, however, axisymmetric instabilities can be excited with a much lower toroidal field of about 5,000 Gauss. A few recent calculations have been done to investigate axisymmetric instabilities in solar tachocline \\citep{cdg2008, dgcm2009, hc2009}. These studies show that latitudinal differential rotation of up to 18\\% pole-to-equator amplitude will be hydrodynamically stable in 2D, but is unstable in 3D or in the presence of magnetic fields.  Various observations, including light curves of rotating starspots \\citep{mg2003}, spectroscopic measurements \\citep{rs2003}, long term changes in Ca II H and K fluxes \\citep{dsb1996}, and Doppler imaging \\citep{bcjd2000} indicate the existence of differential rotation in many stars. Among the stellar population of F, G and K stars, solar-like positive pole-to-equator differential rotation, as well as antisolar pattern in which the equator rotates slower than the pole, are observed. For example, multiwavelength studies of G0-G5V stars by \\citet{mg2003}, and Zeeman-Doppler imaging by \\citet{jd2009} have revealed the existence of solar-like (positive) and antisolar (negative) pole-to-equator differential rotation.  For the spot-dominated late type stars, inversions of light curves for images of dark starspots on the surface reveal the differential rotation pattern of K-stars; so far all of them were found to have solar-like profiles \\citep{rhvh2011}. Antisolar differential rotation has been found among F and G stars, but occurs much less frequently than solar-like differential rotation. Furthermore, the amplitude of antisolar differential rotation in those stars are also found to be smaller than solar type pole-to-equator differential rotation.  Theoretical modeling has enabled us to understand how the solar-like positive pole-to-equator differential rotation is formed in the solar and in stellar convection zones \\citep{gm1986,rempel2005}. Stars with deeper convection zones are likely to have broader differential rotations with latitude \\citep{gilman1979}. Except for anomalously weak rotators, such stars are likely to have equatorial acceleration like the Sun does. On the other hand, stars with shallow convection zones, like F stars, may not have broad differential rotation, but rather a much more structured pattern. Even if they rotate much faster than the Sun, if the turnover time in their shallow convection zone is short enough, they could have antisolar differential rotations.  Hydrodynamic and magnetohydrodynamic instability of solar-like differential rotation in stars has also been explored \\citep{ks1982,uss1996, spruit1999, spruit2002, braithwaite2006}. But to our knowledge for the case of antisolar differential rotation neither detailed theory of formation nor the stability analyses have been carried out yet.  The aim of this paper is to analyze the instability of antisolar differential rotation in order to answer the following questions. (i) Why have only about a dozen antisolar stars been found? (ii) Is antisolar differential rotation more unstable than the solar-like pattern, and could that be the reason? (iii) What is the limiting amplitude of a negative pole-to-equator differential rotation that can be stable? (iv) Why has no antisolar K-star been found yet?  The observations of stellar rotation are made at the surface of a star, and we do not know the differential rotation pattern in their tachoclines. Using the analogy that the Sun's positive pole-to-equator surface latitudinal differential rotation persists down to the tachocline, we assume that the stellar tachoclines reflect the same sign of the pole-to-equator latitudinal differential rotation as observed at their surfaces. The differential rotation formed in the stellar convection zone imposes an upper boundary condition on the tachocline and provides one source for the energy needed to drive instabilities there.  But the precise relation between stellar rotation at the surface and in the stellar tachocline may be more complex, and depend on such factors as the thickness of the stellar convection zone. In the case of the Sun, and probably K stars, the inertia of the convection zone is so large compared to the tachocline below it that it is unlikely that instability in the tachoclines of such stars could significantly change the surface differential rotation pattern. \\citet{rempel2005} did show that thermal effects in the solar tachocline could modify significantly the rotation contours of the convection zone, moving them away from being constant on cylinders as global convection theory tends to create. But with the thin, low density convection zones of F stars, instability in the tachocline could have a much stronger feedback on the surface rotation, resulting in a lower overall differential rotation for F-stars, including at their photospheres. This is a reason to study tachocline instabilities in stars as one determinant of their surface differential rotations.  We will use for our calculations a 3D thin-shell model of the solar/stellar tachocline to analyze axisymmetric and nonaxisymmetric MHD instability of the pole-to-equator latitudinal differential rotation for a wide range of positive and negative amplitudes. We do the MHD instability problem as opposed to the purely HD problem because it is likely that most stellar tachoclines contain substantial toroidal fields. Also, in the previous calculations on magneto-shear instabilities by us and many others, as referred earlier in this section, major findings in 2D, quasi-3D and 3D models were that the magneto-shear instability exists for a wide range of toroidal field bands of latitudinal widths of above $2.7^{\\circ}$, located at a wide range of latitudes above $\\sim 13^{\\circ}$, and for a wide range of peak field strength from a few hundred Gauss and above.  Magneto-shear instabilities occur for differential rotation amplitudes that are characteristic of both radiative and overshoot tachoclines. In the case of the Sun, surface observations can give some guidance about the bandwidth of the tachocline toroidal field bands to be of $6 - 10^{\\circ}$, if active regions are being produced from them (see \\citet{zwaan1978}). Rising flux tube studies (see, for example, \\citet{ffd1993}) indicate that these bands also perhaps spend their maximum time in the latitude range of $0^{\\circ} - 45^{\\circ}$. So, in order to pick a typical case for the Sun and solar-like stars for investigating magneto-shear instabilities with solar and antisolar differential rotation, we consider in this study a $10^{\\circ}$ toroidal band located at $30^{\\circ}$ latitude, and explore the details in the parameter space of solar/antisolar differential rotation amplitude and toroidal field strength. ",
        "conclusions": "From analyses of global MHD instabilities in 3D thin-shell models of the solar/stellar tachoclines, we find that solar-like and antisolar type latitudinal differential rotations with coexisting toroidal bands are unstable. Antisolar differential rotation is, in general, more unstable than a solar-like one in all F, G and K stars -- the pole rotating one or two percent faster than the equator can make the latitudinal differential rotation unstable in weakly magnetized overshoot tachoclines, although it requires much stronger magnetic fields to be unstable in the radiative tachoclines in G and K stars.  High effective gravity in the radiative tachoclines largely suppreses the radial motions by the negative magnetic buoyancy and therefore, it is difficult to grow the perturbations. By contrast, low effective gravity allows the radial motions to grow, and more so when the variation in differential rotation in latitude is smaller. Thus antisolar type differential rotation becomes more unstable than solar type ones, because less $J\\times B$ force is needed to overcome the perturbation centrifugal force that tries to stabilize the displacement.  The eigen functions for the disturbances with solar-like and antisolar type differential rotations have similarities in the high $G$ case, but have important differences in the low $G$ case. In the high $G$ case, the velocity and pressure perturbations are in nearly geostrophic balance and horizontal divergences in flow and field are very small. The E-W tilts of the velocity vectors indicate angular momentum transport towards the equator, but this means that the instability tries to destroy the antisolar type gradient in latitudinal differential rotation, whereas it builds the solar type gradient. However, in both cases, the primary energy source for the disturbances is the toroidal field, so for high $G$, solar and antisolar differential rotations require a similar amplitude of toroidal field to become unstable.  For low $G$, by contrast, the eigen functions have more radial motions and fields and associated horizontal divergences. For weak toroidal fields, the hydrodynamics tend to dominate at all latitudes including at the locations of the toroidal bands. The energy for the growth of the disturbances is primarily coming from the differential rotation. The E-W tilts in horizontal velocity vectors for antisolar type differential rotations imply angular momentum transport towards the equator, which takes kinetic energy out of the differential rotation. Disturbances with the same tilts would actually reinforce a solar type differential rotation, damping the disturbances. This is why instability for solar type differential rotations does not occur until a substantially higher toroidal field is included.  For F stars, antisolar latitudinal differential rotation of only a few percent pole-to-equator amplitude is unstable in both the overshoot and radiative tachoclines, because the effective gravity is much smaller in F stars compared to G and K stars. Therefore, antisolar differential rotation is unlikely to be large in F stars. By contrast, in K stars, antisolar differential rotation is stable in the radiative tachocline but unstable in the overshoot tachocline. This implies that we should find more K stars with antisolar differential rotation compared to F stars, but the observations have so far indicated the opposite, namely no K stars and about half a dozen of F stars are found to have antisolar differential rotations.  We might be able to detect more antisolar K stars in future. One important point we should restate again is that the stability analyses we are performing is in the tachoclines of these stars, and stellar differential rotations are generally observed at the surface. The deeper convection zone in the K stars might mean that the surface is not reflecting the tachocline differential rotation pattern, whereas a very thin convection zone in F stars may be relfecting a 1 or 2\\% stable antisolar tachocline differential rotation at the surface. In any case, antisolar stars are found to have a very small amount of positive pole-to-equator differential rotation.  In this paper, we did not attempt to understand what builds a solar-like or an antisolar type differential rotation, but once they are built, we investigated their stability to disturbances with various longitudinal and radial wave numbers. We focussed on a $10^{\\circ}$ toroidal band placed at $30^{\\circ}$ latitude in the tachoclines of these stars, but other latitude locations of the bands and band-widths can be studied for antisolar differential rotations. Full 3D models have been explored by \\citet{asr2007} and \\citet{zls2003} for the stability analysis of the solar-like differential rotation, and antisolar case is yet to be explored."
    },
    "1107/1107.0741_arXiv.txt": {
        "abstract": "The CoGeNT collaboration has reported evidence of an annual modulation in its first fifteen months of data. Here we compare the amplitude and phase of this signal to the modulation observed by the DAMA collaboration, assuming that both arise due to elastically scattering dark matter (DM). We directly map the CoGeNT signal to the DAMA detector without specifying any astrophysical parameters and compare this with the signal measured by DAMA. We also compare with constraints from CDMS II and XENON10. We find that DM of mass 5-14 GeV that couples equally to protons and neutrons is strongly disfavoured. Isospin-violating DM fares better but requires a boosted modulation fraction. ",
        "introduction": "There is strong observational evidence for a large abundance of particle dark matter (DM) in our Universe. Detecting this matter through non-gravitational interactions is a great challenge. A potentially characteristic signal arising from particle DM, however, is the presence of an annual modulation in the event rate at direct detection experiments \\cite{Drukier:1986tm, Freese:1987wu}. These experiments  aim to detect the energy deposited by a nucleus after scattering with DM. The modulation signal arises due to the motion of the Earth relative to the galactic halo and peaks when the Earth is travelling fastest with respect to the halo. For most dark matter halo models, this occurs in late May or early June (see e.g.\\ \\cite{Fornengo:2003fm,Green:2003yh}).  Two experiments have now detected an annual modulation signal that has many features consistent with a signal arising from DM. The DAMA collaboration \\cite{Bernabei:2008yi, Bernabei:2010mq} has observed an annual modulation over thirteen years; first with the DAMA/NaI setup and later with the DAMA/LIBRA upgrade. The large mass of the target material and the long exposure time mean that the statistical nature of the modulation is beyond question. More recently, the CoGeNT collaboration \\cite{Aalseth:2011wp} has presented tentative evidence for an annual modulation after analysing fifteen months of data.  In this paper, we assume both signals arise due to DM and study the consistency of the two signals assuming the DM scatters elastically. In particular, we consider whether the phase and amplitude of the modulations are consistent. Previous studies \\cite{Frandsen:2011ts, Hooper:2011hd, Foot:2011pi, Belli:2011kw, Schwetz:2011xm} have shown that DAMA and CoGeNT can be in agreement for DM of mass $\\sim5$-$14$ GeV that scatters elastically via a spin-independent interaction, albeit with tension from the CDMS \\cite{Ahmed:2010wy}, XENON10 \\cite{Angle:2011th}, XENON100 \\cite{Aprile:2011hi} and SIMPLE \\cite{Felizardo:2011uw} experiments.\\footnote{See \\cite{Collar:2011kf, Collar:2011wq, Collar:2011kr} for a critical discussion of these experiments.}  However, it is well known that interpreting the signals at direct detection experiments is sensitive to many uncertainties, particularly uncertainties in astrophysical parameters (see e.g.\\ \\cite{Fairbairn:2008gz, MarchRussell:2008dy, Kuhlen:2009vh, McCabe:2010zh, Green:2010gw, Lisanti:2010qx, Arina:2011si}). Therefore, we follow the approach of \\cite{Fox:2010bz}, which specifies how to directly map experimental signals from one detector to another, allowing a comparison without specifying any astrophysical parameters. Given the modulation observed at CoGeNT, we calculate the peak day and modulation amplitude expected at DAMA, which we compare with that observed at DAMA. The phase and amplitude will be the same if both modulations arise from DM. We finish by considering constraints from the low threshold analysis of CDMS II \\cite{Ahmed:2010wy, Akerib:2010pv} and XENON10 \\cite{Angle:2011th}. ",
        "conclusions": "We have presented a comparison of the CoGeNT and DAMA modulation signals free from astrophysical uncertainties, having assumed that both modulation signals arise due to elastically scattering DM.  We found that the peak day of the CoGeNT modulation is always earlier than the peak day of the DAMA modulation. The $1 \\sigma$ confidence regions do overlap in some parameter space, but the best fit points typically differ by $\\sim 30$ days. The SHM, which assumes the DM is distributed isotropically, predicts a peak day close to 2nd June. If CoGeNT continue to measure a peak day at the lower edge of the DAMA confidence regions (around mid-May), there will be interesting consequences for DM galactic halo models.  For DM that couples equally to protons and neutrons ($f_n/f_p=1$), we found the measured modulation amplitude and peak day at DAMA is consistent with that expected based on the CoGeNT results. However, the XENON10 S2 analysis excludes the whole mass range in question. Even if we were to ignore the XENON10 analysis, tension remains with the unmodulated CoGeNT rate and the constraint from CDMS II-germanium, which require large modulation fractions. For $m_{\\chi}=12$ GeV, the CoGeNT unmodulated rate excludes the DAMA modulation signal. Moreover, for all masses we consider, the modulation fraction needs to be larger than $\\sim 70\\%$ to be consistent with the low energy analysis of CDMS II-germanium. The SHM predicts $\\sim 9\\%$, which is significantly smaller. Such large deviations from the SHM seem unrealistic. Therefore, based on the constraints from XENON10, CDMS II and the CoGeNT unmodulated rate, we conclude that elastically scattering DM with $f_n/f_p=1$ is unlikely to be the source of the DAMA and CoGeNT modulation signals.  We also considered DM with isospin violating couplings $f_n/f_p=-0.7$. For this choice, the expected modulation amplitude at DAMA, calculated from the CoGeNT measurement, is generically higher than that observed at DAMA. For $m_{\\chi}=12$ GeV, we find that only the $2 \\sigma$ CoGeNT and DAMA regions overlap. However, for $m_{\\chi}=7$ GeV, amplitudes at the lower end of the CoGeNT $90\\%$ region are compatible. Furthermore, the CDMS II germanium and silicon constraints on the modulation fraction are milder, typically requiring modulation fractions $\\sim 15 \\%$. This is still larger than that from the SHM, so if the DAMA and CoGeNT signals do arise from DM, the constraints from CDMS II indicate that there will be interesting consequences for the DM galactic halo model.  It is clear that the evidence for a modulation in the CoGeNT data is still tentative and that much more data is required to definitively confirm a modulation. Fortunately, the CoGeNT-4 upgrade should provide much more data and significantly shrink the CoGeNT best fit regions. With more data, the approach presented here will serve as a useful complementary test on the consistency of DAMA and CoGeNT. In comparison to the usual method of displaying results in the $\\sigma_n$ -- $m_{\\chi}$ plane, this approach has the advantage that the results do not depend on any astrophysical parameters. \\newline  \\noindent{{\\bf Note added:} Following the submission of v.1 of this work to the arXiv, related work appeared: \\cite{Farina:2011pw} and \\cite{Fox:2011px}. Our conclusions are in agreement."
    },
    "1107/1107.2058_arXiv.txt": {
        "abstract": "We present deep {\\it HST, Chandra, VLA} and  {\\it ATCA} images of the jets of PKS  0208--512 and  PKS 1202--262, which were found in a {\\it Chandra} survey of a flux-limited sample of flat-spectrum radio quasars with jets (see Marshall et al., 2005). We discuss in detail their X-ray morphologies and spectra.  We find optical emission from one knot in the jet of PKS 1202--262 and two regions in the jet of  PKS 0208--512.  The X-ray emission of both jets is most consistent with external Comptonization of cosmic microwave background photons by particles within the jet, while the optical emission is most consistent with the synchrotron process.  We model the emission from the jet in this context and discuss implications for jet emission models, including magnetic field and beaming parameters. ",
        "introduction": "Relativistic jets appear to be a common result of accretion onto compact objects.  The jets of active galactic nuclei (AGN) include objects with a wide range of luminosities and sizes.  In the most powerful sources -- quasars and Fanaroff-Riley type II (FR II, Fanaroff \\& Riley 1974)  radio galaxies -- the jets terminate  at the hot spots,  compact bright regions hundreds of kpc away from the nucleus of the host galaxy, where the jet flow collides  with the intergalactic medium  (IGM), inflating the radio lobes.  AGN jets are enormously powerful, with a total bolometric power output that is often comparable to or greater than that from the host galaxy, and a kinetic energy flux that can be comparable to the AGN's bolometric luminosity (Rawlings \\& Saunders 1991).  Until the last decade, almost all progress towards understanding the physics of jets had come either from numerical modeling or multi-frequency radio mapping. However, the pace of discovery has accelerated greatly in the past decade with {\\it HST} and {\\it Chandra} observations.  One of the first observations by {\\it Chandra} is a perfect illustration:  the target, PKS 0637--752, was a bright radio-loud quasar. It was believed that this source would be unresolved in the X-rays, and therefore ideally suited to focus the telescope. Instead, we saw a beautiful X-ray jet well over 10 arcseconds long (Chartas et al. 2000, Schwartz et al. 2000), with morphology similar to that seen in the radio.   Deeper multi-band observations of this jet have since been done to constrain the nature of its emission and also its matter/energy content  (Georganopoulos et al. 2005, Uchiyama et al. 2005, Mehta et al. 2009).  Indeed, since the launch of the two Great Observatories, the number of extended, arcsecond-scale jets known to emit in the optical and X-rays has increased about ten-fold, from less than 5 to more than 50, including members of every luminosity and morphological class of radio-loud AGN.  For jets, multiwaveband observations give physical constraints that cannot be gained in any other way.  In the bands where synchrotron radiation is the likely emission mechanism (for the most part, radio through optical), one can combine images to  derive information regarding particle acceleration, jet orientation, kinematics and dynamics, electron spectrum and other information.  The X-ray emission mechanism can be either synchrotron (with very short particle lifetimes) or inverse-Compton in nature (for recent reviews, see Harris \\& Krawczynski 2006, Worrall 2009).  In either case, the X-ray emission gives a new set of constraints on jet physics, because it represents either (in the case of synchrotron radiation) the most energetic  particles ($\\gamma \\sim 10^6-10^8$) in the jet, which must be accelerated {\\it in situ}, or (in the case of inverse-Compton emission) the very lowest energy ($\\gamma \\sim$ few - hundreds) particles, a population that controls the jet's matter/energy budget yet cannot be probed by any other observations, and which still must be linked to the low-frequency radio emission.  Combining information in multiple bands can then give unique information on the magnetic field and beaming parameters as well as the matter/energy budget of the jet.  Following the discovery of bright X-ray emission from the jet in  PKS 0637--752, we embarked on a survey (Marshall et al. 2005, hereafter Paper I) to assess the rate of occurrence and properties of X-ray jets among the population of quasars.  Our initial goals were to assess the level of detectable X-ray fluxes from radio-bright jets, to locate good targets for detailed imaging and spectral followup studies, and where possible to test models of the X-ray emission by measuring the broad-band, spatially resolved,  spectral energy distributions (SED) of jets from the radio through the optical to the X-ray band.  The survey has been largely successful in achieving these  aims, thanks to a relatively high success rate of about 60\\% of the first 39 observed (Marshall et al., 2011).  We have used the survey data ($\\sim 5$ ks {\\it Chandra} observations) along with other multiwavelength data  (including both longer {\\it Chandra}  observations as well as deep observations in other bands) to study the spectral energy distributions and jet physics of several of the jets in our sample (Gelbord et al. 2005; Schwartz et al. 2006a (hereafter Paper II), 2006b, 2007; Godfrey et al. 2009).   Here we discuss relatively deep (35-55 ks) {\\it Chandra} and {\\it HST} (2--3 orbits) observations of PKS B0208--512 and PKS B1202--262 (hereafter 0208--512 and 1202--262 respectively), two of the X-ray brightest jets found in our survey.  Both had 40-115 counts detected in the initial {\\it Chandra} survey observation,  which detected multiple jet emission regions in the X-rays as well as make an initial assessment of the jet physics (Paper II). The goal of these deeper observations was to examine these conditions in greater detail, assess the jet's morphology and spectral energy distribution in multiple bands, and constrain the X-ray spectrum and emission mechanism.  0208--512 and 1202--262 were first discovered in the Parkes Survey of Radio Sources (Bolton et al. 1964).  0208--512 was first identified by  Peterson et al. (1979) as a quasar at $z=1.003$, and is a powerful X-ray and Gamma-ray emitter.  It was one of the  first extragalactic sources detected by both EGRET (Bertsch et al. 1993) and COMPTEL  (Blom et al. 1995), and has highly unusual gamma-ray properties, being one of the  few blazars characterized by flares at MeV energies, rather than in the GeV or TeV  domains (Stacy et al. 1996, Blom et al. 1996).  1202--262 was identified as a quasar by Wills, Wills \\& Douglas (1973) and Peterson et al. (1979).  It is at a  redshift of $z=0.789$.  In \\S 2 we discuss the observations and data reduction procedures, while in \\S 3, we discuss results from these observations.  In \\S 4, we present a discussion of physical constraints that can be gained from these observations. Finally, in \\S 5, we sum up our conclusions. Henceforth, we use a flat, accelerating cosmology, with $H_0 = 71 {\\rm km ~s^{-1} ~ Mpc^{-1}} $ (for consistency with Paper II), $\\Omega_M=0.27$ and  $\\Omega_\\Lambda=0.73$. ",
        "conclusions": "We have presented new, deep images of the jets of  0208--512 and 1202--262 in three bands:  the radio, optical and X-rays, using multiple telescopes.  These data allow us to construct detailed models of these jets and their emission processes, and comment on physical parameters and how (if) they change with distance from the quasar core.  In addition,  as these sources are among the brightest X-ray jet sources known, the results of such an analysis has implications for our knowledge of the class as a whole, since the number of sources where such an analysis has been done is still small. We have plotted in Figure 11 the luminosity of various jet regions in each source.  Here we proceed to discuss the nature of each of the jet regions in terms of the most popular jet emission models, under which the X-ray emission  seen by {\\it Chandra} is either synchrotron emission or CMB emission that has been Comptonized by the synchrotron-emitting jet particles.  We discuss each of these in turn.  \\subsection{Synchrotron X-ray Emission}  One possibility is that the X-ray emission is due to  synchrotron emission, which is the most likely mechanism in less powerful radio galaxy jets (e.g., Worrall et al. 2001, Perlman \\& Wilson 2005, Harris \\& Krawczynski 2006).  This model was applied to quasar jets by Dermer \\& Atoyan (2002; see also Atoyan \\& Dermer 2004), in the context of a single electron population under which the upturn in the  broadband spectrum would be explained by the fact that the Klein-Nishina limit becomes relevant at these energies, thus imposing a different energy  dependence to their spectral aging.  That model cannot succeed for most of the bright X-ray quasar jets, because it is unable to produce jets where the optical emission is well below the $\\nu F_\\nu$ level seen in the X-rays. As can be seen in Figure 11, both of  our jets are in this regime, with the exception of the optically bright knot $0.4''$ north of the core of 1202--262, which is not clearly resolved.  However, the brightness of this knot's optical emission makes its SED very different from the other knots (Figure 11).  A more elaborate version of the synchrotron model holds that a second, high-energy population of electrons is present (Schwartz et al. 2000, Hardcastle 2006, Harris \\& Krawczynski 2006, Jester et al. 2007).   The X-ray spectral indices we observe in the jet ($\\alpha = \\Gamma - 1 = 0.7$ in the case of the jet of 0208--512 and $\\alpha = 0.5$ in the case of the jet of 1202--262) require electron energy distributions (EEDs) that  locally would be a power law $n(\\gamma) \\propto \\gamma^{-p}$, with $p = 2 \\alpha_x +1 = 2 - 2.4$.  However, if these electrons are in the regime where their radiative aging is not determined by the  Klein-Nishina cross-section, and if the cooling time is significantly faster than their escape time, then the aceleration mechanism that produced them is required to provide an EED of $n(\\gamma) \\propto \\gamma ^{-1}$ to $\\gamma^{-1.4}$. Note that adiabatic losses do not steepen the electron energy distribution and thus result in similar spectral indices (Kirk et al. 2000)  Two-zone synchrotron models effectively double the number of free parameters needed to model the broadband jet emission. Moreover, they force one to consider different and physically distinct acceleration mechanisms and zones for the electron populations producing the emissions seen in X-rays and at lower frequencies.  In particular, the nature  of the optical ``valley\" seen in these jets requires that the X-ray emitting particles must be accelerated at energies of at least $\\sim 10$ TeV and then accelerated up to at least 100 TeV before they escape, necessarily to an environment of much lower magnetic field, so that they do not produce substantial optical-UV synchrotron emission as they cool (see Georganopoulos et al. 2006 for details). Therefore, we do not favor these models, although we cannot rule them out.   \\subsection{Comptonized CMB X-ray Emission (EC/CMB)}  A more commonly invoked model is that the X-ray emission observed from these jets is  due to inverse-Comptonization of the cosmic microwave background (the so-called EC/CMB model, Celotti et al. 2001, Tavecchio et al. 2000). This model has the  virtue of requiring only a single population of electrons, and in addition takes advantage of a mandatory process to produce the high-energy emission.  Central to this model is the requirement that the jet remain relativistic out to  distances of many kiloparsecs from the quasar nucleus with bulk Lorentz factor  $\\Gamma\\sim 10$. This is significantly larger than the constraints form radio  jet to counter jet flux ratios (Arshakian \\& Longair 2004) that require  $\\beta=u/c \\gtrsim 0.6$  ($\\Gamma\\gtrsim 1.25$).  The EC/CMB model requires a continuation of the EED to low Lorentz factors -- energies which currently cannot be observed in any other way.  Electrons at these very low energies by necessity will dominate the overall energy budget of the  jet (just because of their sheer numbers) and thus the models are also constrained by the need to not violate the Eddington luminosity for the supermassive black hole (see  Dermer \\& Atoyan 2004 for the general argument and Mehta et al. 2009 for  its detailed application on PKS  0637--752). To make use of the large $\\Gamma$,  jets have also to be well aligned to the line of sight $\\theta\\sim 1/\\Gamma$  and this poses additional constraints in the model because the actual size of  the jet becomes at least $\\Gamma$ times the projected size (Dermer \\& Atoyan 2004).  In the EC/CMB model the X-ray emission is related to very low-energy electrons and thus we would naively expect to see a very smooth X-ray emission, outwardly similar to what is observed at low radio frequencies.  This is definitely what is observed in the jet of  1202--262, where the X-ray morphology is very smooth, without significant knots (Figure 9), much more similar to  what is seen in the 4.8 GHz VLA map than to what we see at 20 GHz. However it  is hard to comment on whether this is the case in  0208--512 given the smaller angular size of its jet.  Knotty X-ray emission could also be expected in the EC/CMB model under two scenarios.  The first of these is that the injection in the jet varies and that the X-ray knots correspond to periods of  increased plasma injection that are now traveling downstream. Alternately, the well known ``sausage'' instability due to a mismatch in pressure between the jet and external medium could create alternating compressions and rarefactions.  In our previous work (in particular Paper II), we adopted an approach patterned after the work of Felten \\& Morrison (1966), who showed that under the assumption  that the X-ray emission from all the knots arises from inverse-Compton scattering by the same power-law population of electrons which emit the synchrotron  radiation, the ratio of synchrotron to Compton power will be equivalent to the ratio of the energy density of the magnetic field to that of the target photons:  \\begin{equation} {S_{synch} \\over S_{IC} } =  {{(2 \\times 10^4 T)^{(3-p)/2} B_{\\mu G}^{(1+p)/2}} \\over {8 \\pi \\rho}}, \\end{equation}  {\\noindent where $\\rho = \\Gamma^2 \\rho_0 (1+z)^4$ is the apparent energy density of the CMB at a redshift $z$ in a frame moving at Lorentz factor $\\Gamma$, $\\rho_0 = 4.19\\times 10^{-13} {\\rm ~erg ~cm^{-3}}$ is the local CMB energy density, and the apparent temperature of the CMB is $\\delta T$, where $\\delta = [\\Gamma (1-\\beta \\cos \\theta)]^{-1}$ is the usual beaming parameter, and $T=2.728 (1+z)$ K is the CMB temperature corresponding to epoch $z$.   }  This approach makes two assumptions. First, the CMB  energy density in the comoving frame [$(4/3)U_{CMB, local}\\Gamma^2$]  is treated as isotropic (this is stated in the discussion before eq. 3 of Schwartz et al. 2006).  And second, the beaming of synchrotron and EC/CMB is the same and for that reason no $\\delta$ appears in the ratio.  However, neither of these is required by the physics and in fact both are problematic.  \\subsection{Independently Solving for $B, \\delta$}  We can make further progress on understanding the physical conditions in our emitters by making use  of the formalism adopted by Georganopoulos, Kirk \\& Mastichiadis (2001)  for blazar  GeV emission.  We consider a blob of plasma permeated by a magnetic field $B$,  moving relativistically with a bulk Lorentz factor $\\Gamma$ and velocity $u=\\beta c$ at an angle $\\theta$ to the observer's line of sight.  In the frame of the blob, then, the electrons are  characterized by an isotropic power-law density distribution,  \\begin{equation} n^\\prime(\\gamma^\\prime) = {{k}\\over{4\\pi}} \\gamma^{\\prime -p} P(\\gamma_1, \\gamma_2, \\gamma^\\prime), \\end{equation}  {\\noindent where $\\gamma^\\prime$ is the Lorentz factor of the electron (in the blob's frame of reference), $k$ is a constant, and $P(\\gamma_1, \\gamma_2, \\gamma^\\prime)=1$ for $\\gamma_1 \\leq \\gamma \\leq \\gamma_2$ and zero otherwise. Under the assumption that $\\gamma^\\prime >>\\Gamma$, one can treat the electrons as a photon gas, so that we have a radio spectral index $\\alpha=(p-1)/2$.  If we make use of the Lorentz invariant quantity $n/\\gamma^2$, the Lorentz factor $\\gamma$ of an electron in the lab frame is then $\\gamma=\\delta \\gamma^\\prime$, where the Doppler factor $\\delta=[\\gamma(1-\\beta\\cos\\theta)]^{-1}$.  Then, the electron density $n(\\gamma)$ in the lab frame is }  \\begin{equation} n(\\gamma,\\mu) = n^\\prime(\\gamma^\\prime)\\Big({{\\gamma}\\over{\\gamma^\\prime}}\\Big)^2 = {{k}\\over{4\\pi}} \\delta^{2+p}\\gamma^{-p} P(\\gamma_1\\delta,\\gamma_2\\delta,\\gamma), \\end{equation}  {\\noindent where $\\mu=\\cos\\theta$. Given that the effective volume $V_{\\rm eff}$ of the blob in the lab frame is $V_{\\rm eff}=V~\\delta$  -- as shown in Georganopoulos et al. (2001) -- with $V$ being the blob's volume in its own frame, the energy distribution of the effective number of electrons $N_{\\rm eff}(\\gamma,\\mu)$ is }  \\begin{equation} N_{\\rm eff}(\\gamma,\\mu) = n(\\gamma,\\mu) V_{\\rm eff} = {{kV}\\over{4\\pi}} \\delta^{3+p}\\gamma^{-p} P(\\gamma_1\\delta,\\gamma_2 \\delta,\\gamma). \\end{equation}  Using this,  we obtain for the external Compton luminosity:  \\begin{equation} L_{IC}(\\nu)=c_2k\\delta^{4+2\\alpha} \\nu^{-\\alpha} \\end{equation}  {\\noindent for a moving blob.  Similarly, for a standing feature through which the plasma flows relativistically we would obtain}  \\begin{equation} L_{IC}(\\nu)=c_2k\\delta^{3+2\\alpha} \\nu^{-\\alpha}, \\end{equation}  {\\noindent where the constant $c_2$ can be found in Mehta et al. (2009). Note that this equation  does not depend on the bulk Lorentz factor, but it does assume $\\gamma_2>>\\gamma_1$. The next equation comes from the observed synchrotron luminosity:}  \\begin{equation} L_s(\\nu)=c_1 k \\delta^{3+\\alpha} B^{1+\\alpha}\\nu^{-\\alpha} \\end{equation}  {\\noindent for a moving blob or}  \\begin{equation} L_s(\\nu)=c_1 k \\delta^{2+\\alpha} B^{1+\\alpha}\\nu^{-\\alpha} \\end{equation}  {\\noindent for a standing  feature through which plasma flows relativistically, and  constant $c_1$ can be found in Mehta et al. (2009).  We now have two equations and three unknowns, $B$, $k$, and $\\delta$. Note that the bulk motion Lorentz factor does not enter our equations. }   To close this system of equations we require that the source is in the equipartition configuration. Following Worrall \\& Birkinshaw (2006) the magnetic field $B$ in equipartition conditions is related to the electron normalization $k$ through \\begin{equation} {B^2 \\over 8\\pi}={\\alpha+1 \\over 2} (\\chi+1){k m_e c^2 (\\gamma_1^{1-2\\alpha}-\\gamma_2^{1-2\\alpha})\\over V}, \\end{equation}  {\\noindent where $\\chi$ is the ratio of cold to radiating particle energy density, $V$ is the source volume, and $\\alpha$ is the radio spectral index. If  we make the assumption that $\\alpha>1/2$, as justified by knot observations, and $\\chi=1$, we obtain}  \\begin{equation} {B^2 \\over 8\\pi}=(\\alpha+1){k m_e c^2 \\gamma_1^{1-2\\alpha}\\over V}. \\end{equation}      Now we have three equations (EC/CMB luminosity, Synchrotron luminosity, and equipartition), and three unknowns, $B$, $k$, $\\delta$. These may be solved to uniquely determine them. The values of $B$ and $k$ fix the energy density in the blob. A choice of $\\theta$ provides a bulk Lorentz factor $\\Gamma$ and vice versa.  With a choice of  $\\Gamma$ in hand, we can then calculate the jet power.  Note that the above equations give a somewhat different approach than our previous work.  The differences  lie in the beaming parameters, which come from the different expressions we have for the EC/CMB emission. Both approaches agree on the equations for the synchrotron luminosity (eqs. (7) and (8)).  If one adopts the Georganopoulos et al. (2001) formalism that we use here,  then using the results of equations (7) and (9) we obtain the following for the ratio of the two luminosities $L_s$ and $L_{IC}(\\nu)$:  \\begin{equation} {L_{s}(\\nu)\\over L_{IC}(\\nu)}\\propto \\delta^{-(1+\\alpha)} . \\end{equation}  Note that this ratio is independent of bulk $\\Gamma$ or $\\mu$.  Similar results were derived by Dermer (1995), with an extra multiplication factor of $[(1 + \\beta)/(1+\\mu)]^{(1+\\alpha)}$, which, as discussed in Georganopoulos et al. (2001), comes from the approximation that the seed photons in the frame of the blob for inverse Compton scattering are coming from a direction opposite to the direction of the blob velocity.  The Dermer (1995) equations are presented in a form independent of the system of units by Worrall (2009).  The extra multiplication factor is so close to unity for bulk speeds and angles to the line of sight appropriate for quasar jets (and in particular for the sources presented here), as to render the formalisms essentially identical.  Dermer \\& Atoyan (2004) and Jorstad \\&  Marscher (2004) also used very similar formalisms in their analysis. The fundamental advantage to this approach is that it allows us to calculate $B$ and $\\delta$ directly, with the Lorentz factor of the jet $\\Gamma$ then  resulting because of the choice of angle $\\theta$.  This allows us to find directly a  maximum viewing angle for the jet, $\\theta_{max}$ (i.e., where $\\delta = 2 \\Gamma$).  \\subsection{Application to our Data}  \\begin{figure*}[t]  \\plottwo{0208_plotit.ps}{1202_plotit.ps}  \\caption[]{Spectral energy distributions (SEDs) for the core plus four regions of the jet of  0208--512 (left) and  1202--262 (right).  For the jet of 0208--512, data points for K0 are shown as diamonds, those for K1 are triangles, those for K2 are squares, and the data for K3 are X's.  For the jet of 1202--262, data points for R0 are diamonds, data for R1 are triangles, those for R2 are squares, and the data for R3 are X's.  With the exception of R0 in the 1202--262 jet, all the F475W points are upper limits, as is the case for the F814W points for 0208--512 K2, 1202--262 R1, 1202--262 R2 and 1202--262 R3.  The most likely interpretation for the optical emission of the detected knots of both jets is the extreme tail of synchrotron emission, while the X-ray emission is most likely due to inverse-Comptonization of CMB photons (IC-CMB), as discussed in \\S 4.2. We have overplotted Synchrotron/IC-CMB models of jet components, as detailed in \\S 4.2.2.  For 0208--512, we show models for regions K0 (red), K1 (blue) and K3 (green); a model is not shown for K2 since it is essentially identical to that for K2.  For 1202, we show only the model for R1, as the models for R2 and R3 are essentially identical.  No model is fitted for R0 in the 1202 -262 jet as the knot is not fully resolved in X-rays, making it difficult to constrain the IC-CMB component.}  \\end{figure*}  We now apply the method laid out in \\S 4.3 to calculate physical parameters for the jets of 0208--512 and  1202--262, under the assumption that the X-ray emission from them arises via the EC/CMB scenario. We adopt a single spectral index $\\alpha = (p-1)/2$, equal to that observed in the radio  (i.e., 0.69 for 0208--512 and 0.52 for 1202--262).    These spectral indices are consistent  with those observed in the X-rays, as would be expected under the EC/CMB model. We adopt a uniform magnetic field and assume the equipartition condition.  We also assume that the emitting volume is a cylinder, with a filling factor of unity and assume that the ratio of proton to electron energy  density is also 1. All emitting  volumes are taken to be cylinders with the measured angular lengths, and radii of 0.5 arcseconds.  We assume a moving blob for each feature with the given values of $\\delta$, $B$, etc. (as opposed to a standing shock). The minimum particle energy was chosen as $\\gamma_{min}=15$. These assumptions result in the values laid out in Table 3. The resulting spectral energy distributions are plotted in Figure 11.  The typical error ranges on these parameters are $\\sim 10\\%$, so that, for example, the jet of PKS 1202--262 is consistent with a constant $\\delta$ and $\\theta$ for all components.  Note too that in estimating these parameters we assumed that the same volume is radiating for both the radio and X-ray, whereas in Figure 9 and accompanying text we have noted differences between the X-ray and radio morphologies for both jets. To calculate the jet kinetic power, one needs to select a viewing angle.  For PKS 0208--512, we have chosen a value of $\\theta = 3^\\circ$ , which was the mean angle to the line of  sight for the MOJAVE sample (Hovatta et al. 2009), whereas for PKS 1202--262, which is required to be more highly beamed  ($\\delta>20$ and $\\theta < 2.5^\\circ$) we chose $\\theta=2^\\circ$.  The jet powers reported in Table 3, would increase by a factor of $m_p(p-2)/[m_e(p-1)\\gamma_{min}]$ if we assume instead that for every radiating electron there is a cold proton. The result would be an increase in jet power by a factor of $\\approx 34$ for 0208--512 and $\\approx 2$ for 1202--262. Even with this increase -- which corresponds to the limiting case, under rough equipartition conditions, of the absence of positrons -- the jet powers remain smaller or at most comparable to the Eddington luminosity of a $10^9 M_\\odot$ black hole. \\begin{deluxetable}{ccccccc}    % \\tablewidth{0pt} \\tablecaption{Properties of the X-Ray Jet Components } \\tablecolumns{7} \\tablehead{ \\colhead{Object}& \\colhead{Region} & \\colhead{$B ~(\\mu $G)} & \\colhead{$\\delta $} &\\colhead{$\\theta_{\\rm max}$ (deg)}  & \\colhead {$L_{com}^{\\rm a}$}  & \\colhead{$P_{jet}^{\\rm b}$}}   \\startdata  0208--512 & K0 & 8.7 & 10.2    & 5.6 & 1.03  & 2.10 % \\\\ &\tK1 & 8.8 & 8.3  & 6.9 & 1.06  & 1.36  % \\\\ &     K2 & 12.0 & 7.5  & 7.7 & 1.96 & 2.00  % \\\\ &     K3 & 13.9 & 9.5   & 6.0 & 2.65 & 4.60 % \\\\ 1202--262 & R1 & 4.9   & 22.5 &2.5 & 0.31 & 3.84 % \\\\ &    R2 & 4.2   & 24.8 & 2.3 & 0.23 & 3.99 % \\\\ &   R3  & 5.3   & 23.4 & 2.4 & 0.36 & 1.42 %  \\enddata  \\tablenotetext{a}{Comoving luminosity, in units of $10^{44} ~ {\\rm erg ~s^{-1}}$.} \\tablenotetext{b} {Kinetic power, in units of $10^{45} {\\rm ~erg ~s^{-1}}$.} \\end{deluxetable}"
    },
    "1107/1107.0442.txt": {
        "abstract": "{Recently Block published an astro-ph{\\footnote{http://arxiv.org/abs/1106.3928 (2011).}} insinuating that Lemaitre discovery paper of the Expanding Universe was censored   prior to its translation into English and publication in the Monthly Notices of the Royal Astronomical Society. Consequently, Lemaitre's credit for the discovery of the velocity-distance correlation was not recognized. We examine here the chain of events leading to the discovery of the 'Hubble law'. Our summary: (a) Lemaitre found a theoretical linear correlation between velocity and distance. (b) Lemaitre assumed the existence of a linear relation between velocity and distance and calculated the coefficient. (c) Hubble took the data plotted it and demonstrated that a linear relation represents the observed data.  (d) Hubble never believed in Lemaitre's solution, namely in an expanding universe. Consequently, Hubble never cited Lemaitre.   We conclude that the charge that Lemaitre's  paper was censored or ignored let alone plagiarized by Hubble, is not founded , and explain why Lemaitre's earlier theoretical discovery and derived 'Hubble constant' was not cited or  recognized, by Hubble as well as by many other leading researchers. ",
        "introduction": " ",
        "conclusions": "In table \\ref{tab:early-H} we give chronological order and the values of the early measurements of the Hubble constant. \\begin{table}[htdp] \\caption{Early measurements of the 'Hubble constant'} \\begin{center} \\begin{tabular}{lllll} Author & year & value & Number & comments\\cr &          &          & of objects  &   \\cr Wirtz & 1922-24 & 840 &\\cr Lundmark & 1925 & 71 & 43 & Calculated on the basis of his data \\cr Lemaitre & 1927 & 625  & 42 & Weighted \\cr Lemaitre & 1927 & 670  & 42 & unweighted \\cr Hubble & 1929 & 530 & 24 \\cr de Sitter & 1930 & 460 & 54 \\cr Eddington & 1930 & 522 &\\cr Lundmark & 1930 &   604  & 25 &Spiral-formed PASP, {\\bf 42}, 23, 38, (1930)\\cr Lundmark   & 1930   & 876& 17 &Elliptical  PASP, {\\bf 42}, 23, 38, (1930)  \\cr Oort & 1931 & 290 & 44 \\cr Hubble \\& Humason & 1931 & 558 & 68 & Special treatment for clusters of galaxies \\cr \\end{tabular} \\end{center} \\label{tab:early-H} \\end{table}%"
    },
    "1107/1107.1236_arXiv.txt": {
        "abstract": "Many previous studies have determined that the long lasting emission at X-ray, optical and radio wavelengths from gamma-ray bursts (GRBs), called the afterglow, is likely produced by the external forward shock model. In this model, the GRB jet interacts with the circum-stellar medium and drives a shock that heats the medium, which radiates via synchrotron emission. In this work, we carried out a detailed analysis of the late time afterglow data of GRB 090902B using a very careful accounting of the Inverse Compton losses. We find that in the context of the external forward shock model, the only viable option to explain the X-ray and optical data of GRB 090920B is to have the electron energy distribution deviate from a power-law shape and exhibit some slight curvature immediately downstream of the shock front (we explored other models that rely on a single power-law assumption, but they all fail to explain the observations). We find the fraction of the energy of shocked plasma in magnetic field to be $\\sim 10^{-6}$ using late time afterglow data, which is consistent with the value obtained using early gamma-ray data. Studies like the present one might be able to provide a link between GRB afterglow modeling and numerical simulations of particle acceleration in collisionless shocks. We also provide detailed calculations for the early ($\\lae 10^3$ s) high energy ($> 100$ MeV) emission and confirm that it is consistent with origin in the external forward shock.  We investigated the possibility that the $\\sim 10$ keV excess observed in the spectrum during the prompt phase also has its origin in the external shock and found the answer to be negative. ",
        "introduction": "The external forward shock model (see, e.g., Rees \\& M\\'esz\\'aros 1992; M\\'esz\\'aros \\& Rees 1993, 1997; Paczy\\'nski \\& Rhoads 1993; Wijers, Rees \\& M\\'esz\\'aros 1997; Sari, Piran \\& Narayan 1998; Dermer \\& Mitman 1999) has proven to be a very useful concept in the study of GRBs (for a review of this phenomenon see, e.g., Piran 2004).  The relativistic GRB jet or outflow interacts with the surrounding medium of the progenitor star (circumstellar medium or CSM) and drives a forward shock that accelerates the particles in the CSM, which radiate via synchrotron and Inverse Compton mechanisms.  By modeling GRB afterglows one can learn more about the CSM medium properties (its density), general properties of the outflow (total kinetic energy in the shocked medium) and some details of the shock process (fractions of total energy in the shocked fluid imparted to electrons and magnetic fields). Particles are likely accelerated in collisionless shocks by the Fermi process and the resultant electron energy distribution is expected to be a decaying power-law in electron energy with index $p$ (see, e.g., Krymskii 1977, Axford, Leer \\& Skadron 1978; Bell 1978; Blandford \\& Ostriker 1978; Blandford \\& Eichler 1987; Gallant \\& Achterberg 1999; Achterberg et al. 2001; Sironi \\& Spitkovsky 2011).  Various studies have attempted to calculate the value of $p$ and to test for its universality among all bursts (see, e.g., Shen, Kumar \\& Robinson 2006).  Even though in the majority of previous studies a single value of $p$ is assumed for a particular burst, there is the possibility that the distribution function might deviate from a single power-law (see, e.g., Li \\& Chevalier 2001).  In this paper we study the late ($\\gae$ 0.5 d) afterglow of GRB 090902B in the context of the external forward shock model. In Section 2, we study the afterglow data for this GRB.  In Section 3, we present alternatives to the scenario proposed in Section 2. In particular, we devote most of this section on a very detailed calculation of Inverse Compton losses, which might mitigate some of the issues presented in Section 2.  In Section 4, we show that the only viable solution seems to be that the electron energy distribution is a little steeper for the X-ray radiating electrons than for the optical electrons in the external forward shock, that is, there is some curvature in the electron energy distribution spectrum.  In Section 5, we use the results of the previous section to study the early ($\\sim 50$ s) high-energy gamma-ray data of GRB 090902B and find it consistent with the same origin as the late time afterglow emission. In Section 6, we explore the possibility that the additional power-law spectral component found at $\\sim 10$ s, in addition to the prompt Band function, has also an external forward shock origin. We discuss our results in Section 7 and present our Conclusions in Section 8. ",
        "conclusions": "We have analyzed the late time afterglow X-ray, optical and radio data for GRB 090902B in the context of the synchrotron radiation mechanism in the external forward shock model.  We find that a curvature in the power-law electron energy distribution is needed in order to provide a good fit to the late time optical and X-ray data; radiation losses, varying microphysical parameters and an increase in Inverse Compton losses all fail to explain the observed data.  The late time afterglow fit gives an $\\epsilon_B$ (fraction of total energy in the shock imparted to magnetic fields) consistent with shock-compressed circumstellar medium magnetic field of $\\lae 10$ $\\mu$G and $\\lae 60$ $\\mu$G if we take the efficiency for producing gamma-rays to be $\\sim 50\\%$ and $\\sim 80\\%$, respectively.  Particle acceleration in the external forward shock allows us to set a lower limit on the upstream circumstellar medium magnetic field. We find that the field strength in the unshocked medium in the vicinity of GRB 090902B must be at least 2 $\\mu$G in order to produce 100 MeV photons at 50 s$\\lae t \\lae 10^3$ s (Barniol Duran \\& Kumar 2011).  The calculation presented in this paper represents an improvement in our previous afterglow modeling due to a more precise calculation of Inverse Compton losses.  Here, we include Klein-Nishina effects and also relativistic corrections of the outgoing energy of the Inverse Compton scattered photons.  We also calculate the electron energy distribution self-consistently by determining the synchrotron emission and using it to determine the Inverse Compton losses, which in turn modify the electron energy distribution.  The flux calculated at 100 MeV at 50 s using the external forward shock parameters obtained from the late afterglow data is consistent with the {\\it Fermi}/LAT data, confirming our previous claims (Kumar \\& Barniol Duran 2009, 2010). We also calculated the expected external forward shock at 100 keV and 50 s.  At this time, the observed 100 keV light curve is undergoing a fast decay ($\\sim t^{-3}$).  We find that the external forward shock at 100 keV and 50 s is smaller than the observed value by a factor of $\\sim 40$, easily allowing the observed light curve to decay quickly (confirming our earlier results in Kumar \\& Barniol Duran 2010).  We speculate that for some small fraction of GRBs this steep decay will not be seen, instead, one will see a smooth slowly decaying ($\\sim t^{-1}$) light curve emerge after the main prompt, variable, emission is over.  The origin of the high-energy emission at $<50$ s (when the prompt GBM phase is active) is still a subject of debate. It is claimed that GRB 090902B exhibits variability in the LAT band (Abdo et al. 2009), which would make the external forward shock origin for $<50$ s difficult, since this model in its simplest version predicts smooth light curves. However, the fact that for this GRB the $>50$ s 100 MeV light curve connects smoothly -- as a single power-law -- with the $<50$ s one does argue in favor of the external forward shock origin (see, however, Maxham, Zhang \\& Zhang 2010). The delay on the detection of the first 100 MeV photons with respect of the first GBM photons from this GRB can also be explained by the external forward shock model (Barniol Duran \\& Kumar 2011). A dedicated study of the variability during the prompt phase in the LAT band would certainly shed more light on whether the prompt ($<50$ s) 100 MeV emission of this GRB can be originated from the external forward shock.   We find that the cooling frequency at $\\sim 50$ s is $\\sim 100$ MeV for GRB 090902B.  If this is correct, one should be able to track the cooling frequency as it passes through the LAT band at earlier times ($\\nu_c \\propto t^{-1/2}$).  A presence of this behavior in the data of this and other LAT detected GRBs would help confirm an external shock origin for the LAT emission.  Finally, the prompt spectrum of GRB 090902B contains a power-law component in addition to the usual Band function.  We argue against the external shock origin of the extra power-law component at 7 s (Section 6)."
    },
    "1107/1107.4139_arXiv.txt": {
        "abstract": "We present a detailed chemical abundance study of evolved stars in 10 open clusters based on Hydra multi-object echelle spectra obtained with the WIYN 3.5m telescope.  From an analysis of both equivalent widths and spectrum synthesis, abundances have been determined for the elements Fe, Na, O, Mg, Si, Ca, Ti, Ni, Zr, and for two of the 10 clusters, Al and Cr.  To our knowledge, this is the first detailed abundance analysis for clusters NGC 1245, NGC 2194, NGC 2355 and NGC 2425.  These 10 clusters were selected for analysis because they span a Galactocentric distance range \\rgc$\\sim$9--13 kpc, the approximate location of the transition between the inner and outer disk.  Combined with cluster samples from our previous work and those of other studies in the literature, we explore abundance trends as a function of cluster \\rgc, age, and [Fe/H].  As found previously by us and other studies, the [Fe/H] distribution appears to decrease with increasing \\rgc\\ to a distance of $\\sim$12 kpc, and then flattens to a roughly constant value in the outer disk.  Cluster average element [X/Fe] ratios appear to be independent of \\rgc, although the picture for [O/Fe] is more more complicated by a clear trend of [O/Fe] with [Fe/H] and sample incompleteness.  Other than oxygen, no other element [X/Fe] exhibits a clear trend with [Fe/H]; likewise, there does not appear to be any strong correlation between abundance and cluster age.  We divided clusters into different age bins to explore temporal variations in the radial element distributions.  The radial metallicity gradient appears to have flattened slightly as a function of time, as found by other studies.  There is also some indication that the transition from the inner disk metallicity gradient to the $\\sim$constant [Fe/H] distribution of the outer disk occurs at different Galactocentric radii for different age bins.  However, interpretation of the time evolution of radial abundance distributions is complicated by the unequal \\rgc\\ and [Fe/H] ranges spanned by clusters in different age bins. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.1000_arXiv.txt": {
        "abstract": "We study the spindown of isolated neutron stars from initially rapid rotation rates, driven by two factors: (i) gravitational wave emission due to r-modes and (ii) magnetic braking.  In the context of isolated neutron stars, we present the first study including self-consistently the magnetic damping of r-modes in the spin evolution.  We track the spin evolution employing the RNS code, which accounts for the rotating structure of neutron stars for various equations of state.  We find that, despite the strong damping due to the magnetic field, r-modes alter the braking rate from pure magnetic braking for $B\\leq 10^{13}$G. For realistic values of the saturation amplitude $\\alpha_{\\rm sat}$, the r-mode can also decrease the time to reach the threshold central density for quark deconfinement. Within a phenomenological model, we assess the gravitational waveform that would result from r-mode driven spindown of a magnetized neutron star.  To contrast with the persistent signal during the spindown phase, we also present a preliminary estimate of the transient gravitational wave signal from an explosive quark-hadron phase transition, which can be a signal for the deconfinement of quarks inside neutron stars.  \\newpage ",
        "introduction": "Neutron stars are highly compact stars of typical radius $R\\sim 12$ km and mass $M\\sim 1.5M_{\\odot}$ made mostly of degenerate neutron-rich matter at densities up to several times nuclear matter saturation density $\\rho_0=2.5\\times 10^{14}$ g/cc. By tracking their long-term thermal and rotational evolution, we can learn about the nature of matter under the crust. For example, \\citet{Page11} have proposed that the thermal history of the neutron star in Cas A may be indicating the recent onset of neutron superfluidity deep in its interior. Recently, \\citet{Neg11} have shown how rotational evolution is linked to a reorganization of particle composition in the stellar interior, leading to switching on of exotic neutrino emission processes. A neutron star is a complex system with intertwined physical properties that can change on relatively short astrophysical timescales.  In this paper, we will be concerned with the rotational evolution of a highly magnetized isolated neutron star.  The electromagnetic emissions of a neutron star derive from its rotational kinetic energy, and its spindown is usually measured in terms of a braking index, $n$, which is dependent on the magnetic field configuration ($n$=3 for a dipolar field).  Neutron stars can also spindown through gravitational wave emissions associated to the r-mode \\citep{Anders}.  In fact, the observational interest in r-modes comes from the fact that no neutron stars have been found to spin at rates near the maximum allowed frequency (the ``break-up frequency'').  r-modes offer a possible explanation of this fact: in rotating neutron stars, these modes lose energy through gravitational waves, which carry away angular momentum from the star and act as braking radiation.  As the star spins down, its central density increases.  This increase could be sufficient to make baryonic matter undergo phase transitions to more exotic phases of strongly interacting matter, such as quark matter, with possible implications for gamma-ray bursts \\citep{Bez,HZ,Drago2} and the formation of quark stars, if theoretical conjectures about the absolute stability of strange quark matter are realized in nature \\citep{Itoh, Bodmer, Witten}. While it is almost certain that this quark phase, if it exists inside neutron stars, cannot be a free gas of quarks \\citep{Ozel,Weissenborn11}, an interacting phase of quarks still appears to be consistent with the recent finding of a 2$M_{\\odot}$ neutron star \\citep{Demorest}.   The main questions we seek to answer are: given a few initial parameters of the newly born neutron star, such as its spin period, magnetic field and baryonic mass, can we determine which neutron stars are likely to eventually manifest a quark phase in their interior?  If so, how long before the transition to quark matter occurs?  In this work, we take a step towards answering these questions by taking a closer look at the rotational evolution of a newly-born hot neutron star as it spins down, and we focus on two main driving factors - magnetic braking and gravitational radiation from r-modes. We consider different equations of state (EoS) for neutron stars, since the EoS at high density is uncertain and lacking strong empirical constraints \\citep[although a recent statistical analysis shows that the equation of state stiffens at high densities and is consistent with the expected range in certain nuclear parameters;][]{Steiner10}.  Furthermore, second generation gravitational wave detectors such as Advanced LIGO will soon be operational, and our work provides an update for a similar theoretical study \\citep{Lai} performed several years ago with regard to the LIGO detector. We should note here that in contrast to the assumptions in \\citet{Lai}, our work includes the effects of magnetic damping of r-modes in the spindown evolution, which leads to quantitatively different results.  In section \\ref{Theorymag}, we outline the approach and main equations that describe the neutron star's spindown. The theoretical analysis follows in part the works of \\citet{Lai}, as well as \\citet{Drago1} and \\citet{Rezzolla}. Section \\ref{Magresults} collects our main conclusions from this analysis. In section \\ref{GWevo}, we analyze the evolution of the gravitational wave frequency associated to the growing r-mode, along the lines of \\citet{Owen98}, and present results in section \\ref{GWresults}. We conclude in section \\ref{conclude} with a preliminary calculation of the gravitational wave signal from an explosive deconfinement transition (the ``Quark-Nova\"). ",
        "conclusions": "\\label{conclude}  We have studied the role of r-modes in spinning a rapidly rotating and magnetized neutron star down to typical quark deconfinement densities. Apart from the usual spindown associated to magnetic braking, the magnetic damping of r-modes plays an important role in the period evolution of the star.  We find that, for realistic (small) values of the r-mode saturation amplitude $\\alpha_{\\rm sat}$, the time to deconfinement is sped up by the r-mode instability when magnetic fields are of order $10^{12}$G or less. This result reflects the strong dependence of the magnetic damping effect on $B$ and the evolution of $\\alpha$.  Thus, in contrast to the result in \\citep{Lai}, where the r-mode was seen to affect the spindown already for magnetic fields less than $B\\sim 10^{14}$G, the inclusion of the magnetic damping term leads us to conclude that r-mode driven effects on spindown are pronounced only for $B\\lesssim 10^{12}$G.  We also explored the gravitational wave signal generated during the spindown phase as the r-mode first grows then saturates - we follow the signal continuously until the quark deconfinement threshold is reached.  For realistic values of $\\alpha_{\\rm sat}$, the r-mode saturates and leads to a strong signal in gravitational waves, except in case of very soft equations of state and very large magnetic field ($10^{15}$ G or more).  There could be a sizeable fraction of the total rotational kinetic energy emitted as gravitational waves in this epoch, which can last several years, making it detectable in upcoming second and third generation detectors.  Therefore, gravitational waves could be used in this manner to probe the equation of state inside neutron stars.  As shown here, using r-modes is an alternate way in which the EoS can be probed through gravitational waves~\\footnote{We note that recently derived empirical upper bounds on the gravitational power radiated by the Crab pulsar constrain the ellipticity of deformed pulsars \\citep{Pitkin,Abbott,Virgo} with implications for the maximum theoretical elasticity of the neutron star crust \\citep{Horowitz}.  On the other hand, in order to similarly constrain the r-mode amplitude from searches directed at known pulsars, a different range of frequencies and polarizations must be probed \\citep{Owen}, thus the limits placed from the Crab pulsar do not have direct bearing on r-modes.}  Gravitational waves are going to open a new window of observation into our universe. Among the many discoveries that will be made, we anticipate the exciting prospect that neutron star spindown will reveal signs of the elusive r-mode instability as well as signatures for the onset of quark matter in the core. In the appendix we outline a first preliminary estimate of what the gravitational wave signal from a ``Quark-Nova'' would look like. It would be interesting to examine this signal with detailed simulations, especially in cases when the Quark-Nova occurs very shortly after the neutron star is born in a Supernova - the signature of this \"dual-explosion\" in gravitational wave detectors would be two very different signals coming from the same source but separated in time by a few days to few weeks, depending on the time delay between the two explosions. Such a signature would be unmistakeable in upcoming gravitational wave detectors.  \\appendix {\\bf Appendix: Gravitational waves from the quark-hadron phase transition}  It remains an open question as to what happens to the neutron star when the deconfinement density is reached. One possibility is that the entire star is converted to a quark star due to the inherent stability of strange quark matter \\citep{Witten}. If this conversion occurs in an explosive manner, a``Quark-Nova'' could result \\citep{ODD} - what will the gravitational signal from such an event look like? The conversion involves a two-stage process - neutrons in the core dissolve into a $u$ and $d$ quark fluid that is more compact, causing the core to shrink, followed by combustion to $u,d,s$ quarks through leptonic as well as non-leptonic processes. \\citet{Lin} have studied the first stage with Newtonian hydrodynamics and found that quadrupolar and quasi-radial modes are excited during the collapse, leading to gravitational wave emission with an energy output of ~$\\sim 10^{51}$ ergs. However, this work applies to a mixed phase of quarks and nuclear matter and does not consider an explosive transition that begins with non-premixed fluids.  \\vskip 0.2cm  A full numerical treatment of the second stage, an explosive phase transition taking into account fluid motion in 3D is a complex task and beyond the scope of this work.  However, preliminary steps have been taken in this direction.  \\citet{Niebergal10} solved hydrodynamical flow equations for the combustion of neutron matter to strange quark matter in the laminar approximation, including weak equilibrating reactions and strange quark diffusion across the burning front.  The numerical results suggest laminar speeds of $0.002-0.04$ times the speed of light, much faster than previous estimates derived using only a reactive-diffusive description~\\citep{Olinto}.  Turbulent combustion has been addressed in a recent work by \\citet{Herzog} who found that the combustion stops short of converting the entire star to quark matter (the reaction is no longer exothermic).  This was also the conclusion in \\citet{Niebergal10}, though for a different physical reason~\\footnote{\\citet{Niebergal10} found that as the burning front expands and cools, it enters an advection dominated regime, where the upstream (hadronic) fluid velocity advects the interface backwards faster than it can progress due to reactions and diffusion. Consequently, the interface halts short of the neutron star surface.}. Neither of these works continue on to estimate the gravitational signal from the explosive combustion.  Following the hydrodynamical approach of these studies, we have estimated the gravitational wave signal $h(t)$ from the following steps.  \\begin{itemize} \\item Starting from a mechanically stable configuration where quark matter constitutes a small fraction of the stellar core, we initiate combustion at the speeds obtained in \\citet{Niebergal10}. The quark fluid, described by a simple bag model equation of state is subsequently evolved using inviscid hydrodynamical equations for relativistic fluid flow (relativistic Euler equations) coupled with Newtonian gravity~\\footnote{ Although the Poisson equation for gravity violates the speed of light (since it is an elliptic PDE), it is unavoidable unless one uses the full GR.}, assuming an axisymmetric rotating configuration of the star about the $\\hat{z}$-axis. In the cylindrical coordinates $(r, \\phi, \\theta)$, where $\\phi$ is the polar angle and $\\theta$ is the azimuthal angle, the relevant equations are given by $\\partial_{t} (rU) + \\partial_{r}(rF) + \\partial_{z}(rG) = S^{'}$, where \\begin{equation*} U = \\begin{pmatrix} {\\rm D} \\\\ S u_{r} \\\\ S u_{\\theta}\\\\S u_{z}\\\\ \\tau \\end{pmatrix}\\quad F =\\begin{pmatrix} {\\rm D}u_{r} \\\\ S u_{r}^{2} + p \\\\ S u_{\\theta}u_{r}\\\\S u_{z} u_{r}\\\\ (\\tau + p)u_{r} \\end{pmatrix}\\quad G = \\begin{pmatrix} {\\rm D}u_{z} \\\\ S u_{r}u_{z} \\\\ S u_{\\theta} u_{z}\\\\S u_{z}^{2} + p\\\\ (\\tau + p)u_{z} \\end{pmatrix}\\quad S^{'} = \\begin{pmatrix} 0 \\\\ Su_{\\theta}^{2} + p + rDg_{r} \\\\ - Su_{\\theta}u_{r} \\\\ r {\\rm D} g_{z} \\\\ rS(u_{r}g_{r} + u_{z}g_{z}) \\end{pmatrix} \\end{equation*} and where ${\\rm D} = \\rho \\gamma$, $S = {\\rm D} h \\gamma$, and $\\tau = S-{\\rm D}-p$ are introduced solely for the purpose of writing the relativistic Euler equations in a form analagous to the more familiar non-relativistic Euler equations.  In the natural units where $c=1$, the fluid velocity, Lorentz factor, gravitational vector field, and specific enthalpy are respectively given by $\\vec{u}\\equiv (u_r,u_{\\theta},u_z)$, $\\gamma = 1/\\sqrt{1-|\\vec{u}|^{2}}$, $\\vec{g} = -\\nabla \\phi$, and $h = 1 + \\epsilon + p/\\rho$, where $\\epsilon$ is the specific internal energy, $p$ is the pressure, and $\\phi$ is the gravitational potential. In the non-relativistic limit ($\\gamma \\rightarrow 1$), we recover the non-relativistic Euler equations.   \\item The quark matter has a density $\\rho_{q} = 5 \\rho_{n}$ and is described by the Bag model (bag constant $B^{1/4} = 145$ Mev) while the nuclear fluid has density $\\rho_{n} = \\rho_{\\rm sat}=2.5\\times 10^{14}{\\rm g/cm^3}$ and is described by the perfect caloric EoS with an adiabatic index of $1.7$.  At $t$=0, we initiate combustion with a density discontinuity $\\rho_q-\\rho_n$, which is situated at a radial coordinate $r=R/4$ inside the star, and choose an initial burning front speed of $u_r$=0.01$c$, which is in the range of burn velocities obtained in \\citet{Niebergal10}. We then evolve the above equations using a weighted average flux relativistic HLL (Harten, Lax, and Van Leer) solver on a cartesian grid with spatial resolution 0.5 km and a timestep determined by a Courant number of $0.3$. The solver is coupled with the Poisson equation $\\nabla^2\\phi=4\\pi G\\rho$.   \\item Solving for $\\vec{u}(r,z,t)$ and $\\rho(r,z,t)$, we follow \\citet{Zwerger} to compute the quadrupole wave amplitude $A_{20}^{E2}(t)$ for our axisymmetric configuration, and from there the signal $h(t)$ as given by eqn.(21) of \\citep {Zwerger}.   \\begin{equation} h(t)=\\frac{1}{8}\\sqrt{\\frac{15}{\\pi}}{\\rm sin^2\\theta}\\frac{A_{20}^{E2}(t)}{D} \\end{equation}  We assume a source distance $D=20$ Mpc to plot the maximal signal strength $h(t)$ and the corresponding luminosity $L(t)$ in Fig.\\ref{QN-signal}. The origin of the peak in $h(t)$ is a balance between increasing mass outflow and a decreasing density gradient. The luminosity is proportional to the square of $\\dot{h}(t)$ (hence the sharp dip at the maximum of $h(t)$) while the total emitted energy obtained from integrating the luminosity curve is $\\sim 1.4 \\times 10^{48}$ ergs, which is about 5 orders of magnitude smaller than the binding energy of the neutron star ($10^{53}$ ergs) and 2 orders of magnitude less than the energy released in gravitational waves during the first stage of core-collapse, where the energy comes from coupling oscillations to rotational motion~\\citep{Lin,Mar}.  \\end{itemize}  Based on this result, the gravitational wave signal for this stage of the phase conversion would be hard to detect in Advanced LIGO, unless the source is Galactic (within few kpc). Most of the energy released in the phase transition is in the form of latent heat and neutrinos. However, our results are only a preliminary estimate, designed to provide a guide to the expected signal from a Quark-Nova. Simulations of gravitational wave signals from realistic core-collapse supernova models indicate that convective flows driven by neutrino heating \\citep{Muller04} can drive a strong gravitational wave signal.  The typical time for neutrinos to diffuse out of the hot quark star \\citep[$\\approx$ 0.1-1 sec according to][]{Keranen} is about an order of magnitude larger than the duration of the signal due to explosive combustion to strange quark matter, therefore, we can expect neutrino heating to be important in our context as well.  We have not included such convective effects in our present simulation, nor the effect of magnetic fields, which can also modify Rayleigh-Taylor instabilities~\\citep{Lug02}."
    },
    "1107/1107.2832_arXiv.txt": {
        "abstract": "{\\igrA\\ is one of the still unidentified hard X-ray \\inte\\ sources, reported for the first time in the 4th IBIS/ISGRI catalog.} {We investigated the nature of \\igrA\\ by carrying out a multiwavelength analysis of the available archival observations performed in the direction of the source.} {We present first the results of the timing and spectral analysis of all the X-ray observations of \\igrA\\ carried out with \\rosat, \\asca, \\einst, \\swift, and \\xmm,\\ and then use them to search for possible counterparts to the source in the optical, infra-red, radio and $\\gamma$-ray domain.} {Our analysis revealed that \\igrA\\ is comprised of three different X-ray emitting regions: a point-like source, an extended object and a cometary-like ``tail'' ($\\sim$4~arcmin). A possible radio counterpart positionally coincident with the source was also identified.} {Based on these results, we suggest that the emission from \\igrA\\ is generated by a pulsar wind nebula produced by a high-velocity pulsar. \\igrA\\ might be the first of these systems detected with \\inte\\ IBIS/ISGRI.} ",
        "introduction": "\\label{sec:intro}  The IBIS/ISGRI telescope \\citep{ubertini03} onboard \\inte\\ \\citep{winkler03} has been operating since seven years from now, and the latest available catalog of sources detected with this telescope contains more that 700 objects \\citep{bird10}. A large fraction of them belongs to the class of X-ray binaries (26\\%) and active galactic nuclei (35\\%); a relatively small fraction is represented by cataclysmic variables (5\\%), and another 4\\% comprises supernova remnants, globular clusters and soft $\\gamma$-ray repeaters. About 30\\% of the IBIS/ISGRI sources are still unclassified \\citep[see e.g.][for a recent review]{chaty10}.  In this paper, we investigate the nature of the still unidentified \\inte\\ source \\igrA.\\ This source was reported for the first time in the 4th IBIS/ISGRI catalog \\citep{bird10}, and was detected in the IBIS/ISGRI mosaic at a significance level of 5.4~$\\sigma$ (20-100~keV). The best determined \\inte\\ position of \\igrA\\ is at \\mbox{$RA$=165.341~deg,} \\mbox{$Dec$=-61.056~deg} (J2000), with an associated uncertainty of 4.3~arcmin (90\\% c.l.). The averaged source fluxes in the 20-40~keV and 40-100~keV energy band are 0.4$\\pm$0.1~mCrab and 0.6$\\pm$0.2~mCrab, respectively\\footnote{1~mCrab corresponds to 7.57$\\times$10$^{-12}$~\\ferg\\ in the 20-40 keV energy band and to 9.42$\\times$10$^{-12}$~\\ferg\\  in the 40-100~keV energy band.}. The light curve of the source, obtained from the online tool {\\sc Heavens}\\footnote{http://www.isdc.unige.ch/heavens}\\citep{heavens}, does not show evidence of variability. A possible counterpart to \\igrA\\ in the soft X-ray (0.3-10~keV) domain was identified by \\citet{malizia11}. In the \\swift\\,/XRT field of view (FOV) around \\igrA,\\ they detected a source located at \\mbox{$RA$=165.44~deg,} \\mbox{$Dec$=-61.022~deg} (associated uncertainty $\\sim$6''). By using the refined X-ray position, these authors also suggested that \\igrA\\ could be associated to the \\rosat\\ source 1RXH\\,J110146.1-610121 and the serendipitous \\xmm\\ source 2XMM\\,J110147.1-61012.  In Sect.~\\ref{sec:igra} we report on the analysis of all the available X-ray observations performed with \\swift,\\ \\xmm,\\ \\rosat,\\ and \\asca,\\ in the direction of \\igrA.\\ We use the results of this analysis in Sect.~\\ref{sec:counterparts} to search for possible counterparts to the source in the optical, infra-red and radio domains. A discussion on the nature of \\igrA\\ is presented in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper we analyzed all publicly available X-ray observations carried out in direction of the unidentified source \\igrA\\ with \\swift, \\xmm, \\rosat, \\asca\\ and \\einst. The entire \\inte\\ error circle around \\igrA\\ has been covered in the soft (1-10~keV) X-ray energy band with \\swift\\ and \\rosat. The analysis of these data led to the detection of a single X-ray source inside the error circle given by \\inte\\ \\citep[see also][]{malizia11}. However, the deepest observations of the region around \\igrA, performed with the EPIC cameras on-board \\xmm, revealed that the source detected with \\swift/XRT and \\rosat\\ is likely a blend of two distinct objects: a point-like source and an 8.1~arcsec extended object. The same observations revealed also the presence of a cometary-like tail ($\\sim$4~arcmin) extending from the two sources, and comprised within the \\inte\\ error circle around \\igrA. The X-ray spectra of the three objects displayed only marginally significant differences (see Table~\\ref{tab:xmmspec} and  Fig.~\\ref{fig:contours}).\\\\ The \\rosat\\,/HRI and \\asca\\,/GIS observations of the same region were unable to disentangle the three emitting components. This can be ascribed to the narrower energy range coverage of the former instrument and the lower spatial resolution of the telescopes on board \\asca\\ and \\rosat\\ with respect to the EPIC cameras (see fluxes in the 0.2-2.4~keV band in Table~\\ref{tab:xmmspec} and Sect.~\\ref{sec:rosat},~\\ref{sec:asca}). In all cases, we showed that the X-ray emitting properties of the entire region comprised in the \\inte\\ error circle of \\igrA\\ were in agreement among all the available instruments. In Sect.~\\ref{sec:xmm}, we also showed that the joint \\xmm\\/+\\inte\\ spectrum could be reasonably well described using a single absorbed power law model, with a normalization constant between the two instruments $\\simeq 1$. All these results suggest that \\igrA\\ is a persistent X-ray emitter that displayed in the past $\\sim$ 31 years only marginal variations (if any) in the 0.3-100~keV energy band.  The improved X-ray positions of sources N and S permitted also to search for their possible optical and infra-red counterparts. The closest catalogued 2MASS, USNO-A2.0, and USNO~B1.0 objects to sources~N and S lie just outside the 90\\% c.l. error circle determined with \\xmm; further observations with the ACIS telescope on-board \\chan\\ are therefore needed to reduce the likelihood of a chance coincidence. We found a possible radio counterpart close to sources N and S. The radio contours derived from the publicly available MGPS-2 data match reasonably well with those of the X-ray emission detected with \\xmm\\ around sources N and S. The presence of an \\egret\\ source, 3EG~J1102-6103, positionally coincident with \\igrA\\ was also reported.\\\\  Being \\igrA\\ located only 11~arcmin apart from the nearby SNR~\\SNR,\\ we first considered the possibility that the X-ray emission coming from the unidentified \\inte\\ source was related to the SNR. However in none of the investigated wavelengths (X-ray, radio, optical, NIR) we could clearly identify the presence of emitting structures located between \\SNR\\ and \\igrA\\ which could have favored such an association \\citep[see also][]{rosado1996}.  As an alternative and perhaps more viable hypothesis, we considered that the morphology and emission properties of the 4~arcmin extended X-ray tail in \\igrA\\ might resemble the elongated features observed in the case of PSR~B2224+65 \\citep[``the Guitar'' nebula,][]{hui2007x} and PSR~J0357+3205 \\citep{deluca2011}. In both cases a pulsar (PSR) has been firmly detected at one end of the elongated structures, and the cometary-like tails have been tentatively explained under the ``bow-shock pulsar wind nebulae'' (bsPWN) scenario. These elongated cometary-like objects form when a pulsar escapes from the associated supernova remnant and moves with a very high velocity through the interstellar medium \\citep[typically hundreds of km/s, i.e. orders of magnitude higher than the velocities of the cold and warm components of the interstellar medium, see e.g.] [and references therein]{roberts05}. Close to the pulsar, the relativistic wind of the compact object usually gives rise to a sub-luminous cavity in X-rays surrounded by a termination shock. In this region, particles are thermalized and accelerated, thus producing a conspicuous X-ray and radio emission (through synchrotron processes due to the interaction with the local magnetic field). As the pulsar is moving, a bow shock is formed in front of it and the flow of material advects the emitting particles back along the direction of motion, leading to the formation of an elongated cometary structure \\citep[see e.g.][for a detailed discussion]{gaensler2006}. In a few cases, deep X-ray observations with the ACIS telescope on-board \\chan\\ permitted to clearly disentangle the different contributions to the total X-ray and radio emission produced by the structures around a bsPWN \\citep[see e.g. the case of the PWN G359.23-0.82;][]{gaensler2004}.  If the structures detected in \\igrA\\ are indeed originating from the pulsar wind, the X-ray tail can be interpreted as being due to relic electrons spread along the direction of motion of the PSR. Source~N, extended in X-rays and closer to the peak of the radio emission, would represent, according to this interpretation, the compact PWN (with possible contribution coming also from the bow-shock). The detection of a radio source located close to source~N would support this scenario, as radio synchrotron emission is expected from particles accelerated at the termination shock; see the cases of, e.g. PSRs B1853+01 \\citep{frail1996, petre2002}, B1957+20 \\citep{stappers2003}, B1757-24 \\citep{frail1991, kaspi2001}.  While the association between the extended source~N ad the ``tail'' seems morphologically plausible, the association of source~S to this region is less straightforward.  In case source~S is not physically related to the X-ray tail, it might just by chance be located along the line of sight to the extended structure. The pulsar responsible for the formation of the compact and relic PWNe, might, in this case, be unresolved and comprised in the emission detected from source~N.  In case, instead, that source~S, the point-like source, is physically connected to the PWN, it would be reasonable to associate it to the PSR generating the X-ray structures. Under this scenario the misalignment between the direction of the 4~arcmin tail and the axis between sources~N and S would  imply a significant tilt between the proper motion of the PSR and the position of the compact PWN. A similar misalignment, however, was already observed in e.g. PSR~B2224+65, and ascribed to the strong interaction between the fast-moving pulsar and its dense surrounding environment \\citep[][]{hui2007x}. Indeed, as for the case of PSR~B2224+65, also the environment around \\SNR\\ and \\igrA\\ was found to be a complex and dense region, characterized by the presence of several background and foreground molecular clouds \\citep[][and references therein]{filipovic2005}.  Following the joint XMM+ISGRI spectral analysis (see Sect.~\\ref{sec:xmm}), we conservatively assume that both sources~N and S contribute to the ISGRI emission (the tail component has indeed a much lower flux). Spectral properties (hardness and flux stability) similar to those observed in the present case (see Fig.~\\ref{fig:combined}) have already been measured in, e.g. AX~J1838-0655. This source is a young ($\\tau_c \\equiv P/2\\dot{P} = 23$~kyr) pulsar, surrounded by a PWN \\citep{gotthelf2008}. The best-fit power law slope in the 2-100~keV was measured at $\\Gamma=1.5+/-0.2$ \\citep{maliziaAXJ1838.0-0655}. We note however that also the highly-magnetized neutron stars \\citep[``magnetars'', for a review see e.g. ][]{mereghetti2008} share similar spectral properties. SGR~1806-20, for example, showed a power law slope $\\Gamma=1.6-2.0$ in the range 1-100~keV \\citep{esposito2007}. As the 1-10~keV spectra of magnetars are usually (phenomenologically) modelled with a combined black body+power law (BB+PL) model \\citep{mereghetti2008}, the lack of the BB component in the \\xmm\\ spectrum might be an issue for this interpretation. We therefore estimated for source~S (i.e. the component showing the higher 2-10~keV flux) the contribution of an eventual black body spectral component. We used a temperature $kT=0.5$, in agreement with the one measured for SGR~1806-20 \\citep{esposito2007} and for magnetars in general. Due to the relative low S/N, we could only obtain a poorly constrained upper limit on the ratio between the fluxes of the BB and PL components $F_{\\rm BB} / F_{\\rm PL} < 0.1$. This would be compatible with the value measured in the case of SGR~1806-20 \\citep[$F_{\\rm BB} / F_{\\rm PL} =0.03$, ][]{esposito2007}. The present data therefore do not argue against a magnetar hypothesis. We note also that a PWN associated to a magnetar has already been observed, e.g. in the case of AXP~1E1547.0-5408 \\citep{vink2009}.\\\\  Further observations with an X-ray telescope characterized by a finer spatial resolution (i.e. the ACIS on-board \\chan) and increased spectral and timing statistics are required to solve the issues above and firmly establish the real nature of \\igrA. In case these observations will confirm that \\igrA\\ is a newly discovered PWN generated by a high-velocity PSR, this would be the first detection with \\inte\\ of one of these systems (to the best of our knowledge)."
    },
    "1107/1107.2235_arXiv.txt": {
        "abstract": "{Modeling the formation of the ice giants Uranus and Neptune is a long-lasting problem in planetary science. Due to gas-drag, collisional damping, and resonant shepherding, the planetary embryos repel the planetesimals away from their reach and thus they stop growing (Levison et al. 2010). This problem persists independently of whether the accretion took place at the current locations of the ice giants or closer to the Sun.} {Instead of trying to push the runaway/oligarchic growth of planetary embryos up to 10$-$15 Earth masses, we envision the possibility that the planetesimal disk could generate a system of planetary embryos of only 1$-$3 Earth masses. Then we investigate whether these embryos could have collided with each other and grown enough to reach the masses of current Uranus and Neptune.} {We perform several series of numerical simulations. The dynamics of a considered set of embryos is influenced by the presence of Jupiter and Saturn, assumed to be fully formed on non-migrating orbits in 2:3 resonance, and gravitational interactions with the disk of gas.} {Our results point to two major problems. First, there is typically a large difference in mass between the first and the second most massive core formed and retained beyond Saturn. Second, in many simulations the final planetary system has more than two objects beyond Saturn. The growth of a major planet from a system of embryos requires strong damping of eccentricities and inclinations from the disk of gas. But strong damping also favors embryos and cores to find a stable resonant configuration, so that systems with more than two surviving objects are found. In addition to these problems, in order to have substantial mutual accretion among embryos, it is necessary to assume that the surface density of the gas was several times higher than that of the minimum-mass solar nebula. However this contrasts with the common idea that Uranus and Neptune formed in a gas-starving disk, which is suggested by the relatively small amount of hydrogen and helium contained in the atmospheres of these planets.} {Only  one of our simulations ``by chance'' successfully reproduced the structure of the outer Solar System. However, our work has the merit to point out that models of formation of Uranus and Neptune have non-trivial problems, which cannot be ignored and have to be addressed in future work.} ",
        "introduction": "The accretion of Uranus and Neptune is a long-standing problem in planetary science. \\cite{safronov_1969} was the first to point out that the accretion of these two planets from a planetesimal disk at their current locations would have taken implausibly long timescales. This problem was confirmed by \\cite{2001Icar..153..224L} using modern numerical simulations. \\citet{2004ApJ...614..497G,2004ARA&A..42..549G} claimed that the in-situ formation of Uranus and Neptune could have been possible in a planetesimal disk strongly dominated by collisional damping. This claim, however, is not correct because, as showed by \\cite{2007Icar..189..196L}, planetary cores in a disk with strong collisional damping simply open gaps in the planetesimal distribution around their own orbits and stop accreting.  There is now a consolidated view that the giant planets were closer to each other in the past (probably all within $12\\,$AU from the Sun) and that they moved to their current orbits after their formation \\citep{1984Icar...58..109F,1993Natur.365..819M,1995AJ....110..420M, 1999AJ....117.3041H,1999Natur.402..635T,2005Natur.435..459T,2007AJ....134.1790M, 2010ApJ...716.1323B}. Thus, it is no longer necessary to construct a model capable of explaining the formation of Uranus and Neptune at their current, remote locations.  Forming Uranus and Neptune within 12$-$15$\\,$AU from the Sun is in principle easier than forming them at 20$-$30$\\,$AU because the density of solid material was probably higher and the dynamical timescale (i.e. the orbital period) was shorter. However, forming 10$-$15 Earth mass (M$_{\\oplus}$) cores from a planetesimal disk turns out to be difficult at {\\it any} location. In fact, \\cite{2010AJ....139.1297L} showed that, when planetary embryos achieve a mass of 1$-$3$\\,$M$_{\\oplus}$, they tend to scatter the remaining planetesimals away rather than accreting them. With the help of gas-drag, collisional damping and resonant shepherding, the embryos repel the planetesimals away from their reach and thus they stop growing.  In this paper we explore another possible venue for the formation of Uranus and Neptune. Instead of trying to push the runaway/oligarchic growth of planetary embryos up to 10$-$15$\\,$M$_{\\oplus}$, we envision the possibility that the planetesimal disk could generate a system of planetary embryos of only 1$-$3$\\,$M$_{\\oplus}$; then we investigate with numerical simulations whether these embryos could have collided with each other because they converged at specific orbital radii where their radial migration in the gas-disk was stopped by the presence of Jupiter and Saturn. This growth mode by giant collisions would explain in the natural way the large obliquities of Uranus and Neptune.  More specifically, our scenario is based on the consideration that Jupiter and Saturn presumably got caught in their mutual 2:3 mean motion resonance (MMR), which prevented them from migrating further towards the Sun. Instead, after being trapped in the 2:3 MMR, Jupiter and Saturn either migrated outwards or stayed roughly at steady locations \\citep{2001MNRAS.320L..55M,2007Icar..191..158M, 2007AJ....134.1790M,2008A&A...483..633P}. To date, this is the only explanation we have for why our giant planets did not migrate permanently into the inner Solar System.  The presence of Jupiter and Saturn on orbits not migrating towards the Sun would have acted as an obstacle against the inward Type-I migration \\citep{1980ApJ...241..425G} of the planetary embryos from the outer Solar System. More precisely, any planetary embryo migrating towards the Sun would have been, sooner or later, trapped and halted in a mean motion resonance with Saturn. Then, the accumulation of embryos in these resonances could in principle have boosted their mutual accretion. This paper aims at investigating this possibility with numerical simulations.  We are aware of the new result according to which the real migration of planetary embryos is very different from the classical Type-I migration envisioned in ideal, isothermal disks \\citep{2006A&A...459L..17P,2008ApJ...678..483B,2008A&A...485..877P, 2008A&A...487L...9K,2011MNRAS.410..293P}. In particular, in disks with realistic cooling times, migration is expected to be outward in the inner part of the disk and inward in its outer part \\citep{2009A&A...497..869L}. This generates a region in between the inner and the outer parts of the disk where Type-I migration is basically inhibited. Planetary embryos are expected to be in/close to this no-migration zone, which seems to invalidate our assumption that embryos migrated towards the giant planets until they got captured in resonances.  However, \\cite{walsh_2011}, from constraints provided by the terrestrial planet system and the asteroid belt, argued strongly that Jupiter and Saturn migrated outwards over a range of several AUs. What is important for our purposes is the relative motion of Jupiter/Saturn and the embryos. It does not really matter whether Jupiter/Saturn are on fixed orbits and the embryos tend to migrate towards the Sun, or the embryos do not migrate while Jupiter/Saturn move outwards. In fact, in both cases the embryos approach the giant planets until they are captured in MMRs, which may act like a privileged site for embryo clustering and mutual accretion. Thus, in our simulations, for simplicity we assume that Jupiter/Saturn are on non-migrating resonant orbits while the embryos are affected by inward migration, with different migration speeds from one simulation to another. This migration speed can be interpreted as an actual inward migration speed of the embryos (most likely reduced relative to the classical Type-I migration speed in iso-thermal disks), or the outward migration speed of Jupiter/Saturn or a combination of the two.  Two caveats related to our work need to be stated up-front. First, our study assumes that Jupiter and Saturn are fully formed, while the accretion of these planets is by itself an unsolved problem that we do not address here. This may sound strange. However, there is a consensus that Uranus and Neptune formed after Jupiter and Saturn, because they did not accrete nearly as much of gas. This leads to two considerations: (a) Jupiter and Saturn existed already when Uranus and Neptune formed, so that the former should/may have influenced the accretion process of the latter and (b) whatever mechanism allowed the formation of Jupiter (obviously not the presence of a pre-existing giant planet!), it did not work for Uranus and Neptune, otherwise they would have formed nearly at the same time as Jupiter.  The second caveat is that, because our model is based on migration of planets and embryos in a gas-disk, the study should be performed with hydro-dynamical simulations. These simulations, though, are too slow to treat the evolution of tens of bodies for a few millions of years, as we need to perform in this study. The paper by \\cite{2008A&A...478..929M} is, to date, the only attempt to study the accretion of the cores of giant planets using hydro-dynamical simulations, and it shows all the limitations of this technique. Thus, for this study we use N-body simulations, with artificial forces exerted onto the embryos to mimic the migration and tidal damping forces exerted from the disk. This approach also has its own limitations as it does not account for indirect mutual perturbations that the embryos may exert onto each other through the modifications that they induce in the density distribution of the gas-disk.  The structure of this paper is as follows. In Sect.~2, we explain our simulation methods. In Sect.~3 we illustrate some basic ingredients of the dynamics of embryos and giant planets. In particular, we discuss the concept of resonance trapping with Saturn, resonant trapping in mutual embryo-embryo resonances, resonance loading, onset of a global dynamical instability and possible mutual accretion. Just for illustrative purposes, we do this by introducing one embryo at the time at a large distance from Saturn, even though this is NOT how we think the real evolution proceeded. Then, in Sect.~4 we move to more ``realistic'' simulations, where multiple $3\\,$M$_{\\oplus}$ embryos are introduced at the beginning of the simulation over the 10$-$35$\\,$AU range. We test the dependence of the results on the total amount of gas in the disk, the inward migration rate and the total number of embryos. Having realized that several embryos are lost by having close encounters with Jupiter and Saturn which either eject them onto distant orbits or inject them into the inner Solar System, in Sect.~5 we show how the coorbital corotation torque exerted at the edge of Saturn's gap can act like a planet trap \\citep{2006ApJ...642..478M, 2008A&A...483..633P} and prevent large mass losses. In Sect.~6 we discuss how the results depend on the initial mass of the embryos. In Sect.~7 we address the role of turbulence in the disk and Sect.~8 collects the conclusions and considerations that we derive from this study.  We anticipate that our study is not ``successful'', in the sense that our simulations do not typically lead to the formation of only two planets with masses close to those of Uranus and Neptune. However it shows interesting dynamical mechanisms and intriguing consequences that will need to be addressed in detail in future studies, most likely using hydro-dynamical simulations.  Note on nomenclature: hereafter we call embryos the objects of 1-$3\\,$M$_{\\oplus}$ that our simulations start with and core any object that is formed during the simulation by merging at least two embryos. ",
        "conclusions": "The formation of Uranus and Neptune by runaway/oligarchic growth from a disk of planetesimals is very difficult. The accretion rate is not large enough; moreover the accretion stalls when the major bodies achieve a few Earth masses \\citep{safronov_1969,2001Icar..153..224L, 2007Icar..189..196L}.  In this paper, we have explored a mechanism for the formation of Uranus and Neptune that is different from those explored before.  We assume that runaway/oligarchic growth from a planetesimal disk generated a system of embryos of 1$-$3 Earth masses, but not larger, in agreement with the results of \\cite{2010AJ....139.1297L}. Then, we follow the dynamical evolution of such a system of embryos, accounting for their mutual interactions, their interaction with a disk of gas and the presence of fully formed Jupiter and Saturn on resonant, non-migrating orbits. Because of computation speed, we do not perform the simulations with an hydro-dynamical code. Instead, we use an N-body code, and we apply fictitious forces to the embryos that mimic the orbital damping and migration torques that the disk exerts on the embryos. We use two prescriptions for the migration torque: one that ignores the co-orbital corotation torque and one that takes it into account. With the second prescription a ``planet trap'' \\citep{2006ApJ...642..478M} appears at about $10\\,$AU, just inward of the outer edge of the gap opened by Jupiter and Saturn in the disk of gas.  Our results show that the idea works in principle. In many runs, particularly those accounting for a disk a few times more massive than the MMSN and those accounting for the presence of the planet trap, there is significant mass growth. In these cases, the major core beyond Saturn exceeds 10 Earth masses. These results highlight the importance of the planet trap and of the damping effect that the gas-disk has on the orbital eccentricities and inclinations of planetary embryos.  Our simulations typically do not successfully reproduce the structure of the outer Solar System. In the wide range of the end-states due to the stochasticity  of the accretion process, only one simulation is successful, with two cores of $15\\,$M$_\\oplus$ produced and surviving beyond Saturn and no rogue embryos. In general, our results point to at least two major problems. The first is that there is typically a large difference in mass between the first and the second most massive cores, particularly in the simulations showing the most spectacular mass growth. This contrasts with Uranus and Neptune having comparable masses.  The second problem appears mostly if we account for the planet trap, so that more material is kept beyond the orbit of Saturn. The problem is that the final planetary system typically has many more than two embryos/cores. Several original embryos, or partially grown cores, survive at the end on stable, resonant orbits. In several cases bodies are found in coorbital resonance with each other. This is in striking contrast with the outer Solar System, where there are no intermediate-mass planets accompanying Uranus and Neptune. It might be possible that these additional planets have been removed during a late dynamical instability of the planetary system, but the likelihood of this process remains to be proven.  Accounting for some turbulence in the disk alleviates this last problem. However, when turbulence is strong enough to prevent the survival in resonance of many bodies, the typical result is the production of only one core.  In addition to the problems mentioned above, there is another intriguing aspect suggested by the results of our model. We have seen that, in order to have substantial accretion among the embryos, it is necessary that the parameter $f_d$ is large, namely that the surface density of the gas is several time higher than that of the MMSN. However this contrasts with the common idea that Uranus and Neptune formed in a gas-starving disk, which is suggested by the small amount of hydrogen and helium contained in the atmospheres of these planets compared to those of Jupiter and Saturn. How to solve this conundrum is not clear to us.  In summary, our work does not bring solutions to the problem of the origin of Uranus and Neptune. However, it has the merit to point out non-trivial problems that cannot be ignored and have to be addressed in future work."
    },
    "1107/1107.5309_arXiv.txt": {
        "abstract": "%  About 2500 square degrees of sky south of declination -25 degrees and/or near the galactic plane were surveyed for bright outer solar system objects.  This survey is one of the first large scale southern sky and galactic plane surveys to detect dwarf planets and other bright Kuiper Belt objects in the trans-Neptunian region.  The survey was able to obtain a limiting R-band magnitude of 21.6.  In all, 18 outer solar system objects were detected, including Pluto which was detected near the galactic center using optimal image subtraction techniques to remove the high stellar density background.  Fourteen of the detections were previously unknown trans-Neptunian objects, demonstrating that the southern sky had not been well-searched to date for bright outer solar system objects.  Assuming moderate albedos, several of the new discoveries from this survey could be in hydrostatic equilibrium and thus be considered dwarf planets. Combining this survey with previous surveys from the northern hemisphere suggests that the Kuiper Belt is nearly complete to around 21st magnitude in the R-band.  All the main dynamical classes in the Kuiper Belt are occupied by at least one dwarf planet sized object. The 3:2 Neptune resonance, which is the innermost well-populated Neptune resonance, has several large objects while the main outer Neptune resonances such as the 5:3, 7:4, 2:1, and 5:2 do not appear have any large objects.  This indicates that the outer resonances are either significantly depleted in objects relative to the 3:2 resonance or have a significantly different assortment of objects than the 3:2 resonance.  For the largest objects ($H<4.5$ mag), the scattered disk population appears to have a few times more objects than the main Kuiper Belt population, while the Sedna population could be several times more than that of the main Kuiper Belt. ",
        "introduction": "The strong dynamical connection that the trans-Neptunian objects (TNOs) have to the planets makes determining their population and orbital structures valuable for gaining insight into solar system formation and planet evolution.  The Kuiper Belt, a remnant of the original protoplanetary disk, has a ``fossilized'' record of the original solar nebula and subsequent evolution of the solar system. TNOs are likely primitive with significant amounts of volatiles.  The largest TNOs or dwarf planet sized objects are rare but extremely important for several reasons: 1) The brightest few objects are the only ones accessible to high signal to noise spectroscopy techniques that are required to determine surface compositions, such as methane and water ice (Barucci et al. 2008; Trujillo et al. 2011).  These physical characteristics are important in order to understand the formation, origin and composition of the objects and gain insight into planet formation and chemistry in the original solar nebula. 2) The size distribution of the biggest objects in the Kuiper Belt determines if the mass in the Kuiper Belt is dominated by the largest or smallest objects, which is a key metric of planetismal growth scenarios (Kenyon et al. 2008, 2010; Cuzzi et al. 2010).  The size and number of the biggest objects constrain the density and thus planet formation ability of the original solar nebula in the outer solar system.  3) Occultations of stars by the biggest TNOs are possible to predict and observe in order to probe the TNOs sizes, shapes, albedos, and atmospheres (Elliot and Kern 2003; Elliot et al. 2010).  The Palomar 48 inch Schmidt telescope in the northern hemisphere, with one of the largest CCD cameras in the world, was used to survey most of the sky north of -25 degrees declination for the brightest ($m_{R} \\lesssim 21$ mags) TNOs (Trujillo and Brown 2003; Brown et al. 2004, 2005; Brown 2008; Schwamb et al. 2009, 2010).  In these surveys tens of bright TNOs including likely dwarf planets Eris, Makemake, Haumea, Orcus, Quaoar, Sedna and 2007 OR10 were discovered.  These surveys showed that many of the largest Kuiper Belt Objects (KBOs) have relatively large inclinations with the vast majority of KBOs expected to be found within about 20 degrees of the ecliptic (Brown 2008). Extrapolating the Cumulative Luminosity Function (CLF) to the bright end of the KBOs indicates several large KBOs should be discovered in the southernmost parts of the sky that the surveys from the northern telescopes did not image.  The southern hemisphere has not been well-surveyed for distant solar system objects until now because in the past there were no sensitive, wide-field digital imagers on suitable telescopes in the south.  This changed in 2009 when a large wide-field imager was put onto the 1.3 meter Warsaw telescope at Las Campanas in Chile.  The OGLE-Carnegie Kuiper Belt Survey (OCKS) was implemented to search the Kuiper Belt for dwarf planets and bright TNOs through a shallow survey to fainter than 21st magnitude in the R-band from the southern hemisphere.  OCKS covered the area within a few tens of degrees of the ecliptic for declinations less than -25 degrees and the crowded galactic plane fields in the north and south.  Another independent southern sky survey for KBOs was started in late 2009 with the Schmidt telescope at La Silla (Rabinowitz 2010).  This is the first time most of this sky area was searched for outer solar system objects with modern digital CCD detectors. ",
        "conclusions": "\\subsection{Completion Limits of the Kuiper Belt}  Eighteen outer solar system objects were detected in this survey. Fourteen of these objects were new discoveries showing that this region of sky had not been well-searched for bright, distant objects in the past (Tables 1 and 2).  Combining this southern sky and galactic plane survey with the previous large area northern sky surveys (Trujillo and Brown 2003; Brown 2008; Schwamb et al. 2009, 2010) and a recent large Kuiper Belt survey in the south started by Rabinowitz (2010) makes it likely that the Kuiper Belt is now nearly complete to about 21st magnitude in the R-band.  To date, only three areas have not been well searched for bright outer solar system objects: 1) southern fields very distant from the ecliptic ($>20$ degrees ecliptic latitude) and thus unlikely to harbor many bright KBOs, 2) a few hundred square degrees in the northern section of the north galactic plane near the ecliptic and 3) a few hundred square degrees in the northern section of the south galactic plane near the ecliptic (Figure~\\ref{fig:MapOgle}).  There are around 70 known TNOs with apparent magnitudes brighter than 21 in the R-band (see section 4.2), almost all within 20 degrees of the ecliptic. Thus, there is on average one KBO brighter than 21st mag every few hundred square degrees of sky near the ecliptic.  This means there is likely to be only one or two KBOs brighter than 21st magnitude in the few remaining areas yet to be searched.  Though nearly complete to 21st magnitude now, some objects, especially Centaurs and scattered disk objects, have large eccentricities and thus could become brighter than 21st magnitude in the future as they approach perihelion.  The size of an object at the completeness limit depends on the distance and albedo of the object (Figure~\\ref{fig:distancesize21}). The largest few objects, with radii greater than about 500 km, have been found to have very high albedos ($\\rho_{R} \\sim 0.6-0.8$), while smaller objects appear to have moderate albedos ($\\rho_{R} \\sim 0.1-0.2$) (Stansberry et al. 2008).  The high albedos of the largest objects is likely due to atmospheres and/or surface processes such as cryovolcanism (Licandro et al. 2006; Dumas et al. 2007; Rabinowitz et al. 2007; Sheppard 2007).  Assuming a typical albedo of $\\rho_{R}=0.15$ for the moderate sized KBOs of 21st magnitude, the completeness limit at 30 AU is about 80 km in radius, while at 50 AU it is about 225 km in radius (Figure~\\ref{fig:distanceogle}).  In absolute magnitude, H, this would be 6.6 and 4.4 magnitudes respectively (Figure~\\ref{fig:distanceH}).  It is clear that further Pluto or larger sized objects could remain undetected if beyond a few hundred AU.  \\subsection{Size Distribution}  Figure~\\ref{fig:KBOcumH} shows the cumulative number of all known TNOs versus their absolute magnitude, H.  Objects with absolute magnitudes $H > 7$ mags appear to have a roll-over in their size distribution because of detection biases.  The largest KBOs with $H<3$ mags do not follow the simple power-law found for the objects with $3<H<7$ mags (Brown 2008).  The largest KBOs have been found to have preferentially higher albedos, likely because of atmosphere effects and surface activity that keep the surfaces young and bright (Jewitt and Luu 2004; Lykawka and Mukai 2005; Schaller and Brown 2007; Stansberry et al. 2008; Desch et al. 2009).  The absolute magnitudes the largest KBOs would have if they had a more typical albedo of 0.15 (Brown and Trujillo 2004; Brown et al. 2006; Stansberry et al. 2008) are shown by squares in Figure~\\ref{fig:KBOcumH}.  The squares in Figure~\\ref{fig:KBOcumH} fit a simple power-law for all objects with $H<7$ magnitudes ($r < 60$ km assuming 0.15 albedo).  The points in a cumulative distribution are heavily correlated with one another, tending to give excess weight to the faint end of the distribution.  A differential distribution does not suffer from this problem.  Figure~\\ref{fig:KBOdiffH} shows the differential number of all known TNOs versus their absolute magnitude, where, like in Figure~\\ref{fig:KBOcumH}, the largest few objects have had their absolute magnitudes adjusted for their abnormally high albedos compared to smaller objects.  It is clear there is a turnover around an absolute magnitude of 7 mags ($r \\sim 60$ km) showing observational bias beyond this magnitude.  The best fit power-law for the differential points finds $q=3.0\\pm 0.5$ for $H<7$ magnitudes, where $n(r)dr \\propto r^{-q}dr$ is the differential power-law radius distribution with $n(r)dr$ describing the number of TNOs with radii in the range $r$ to $r+dr$.  This is slightly lower than most previous fits ($q\\sim 4$) that were more heavily dependent on fainter (smaller) objects (Jewitt et al. 1998; Trujillo et al. 2001; Petit et al. 2008; Fraser et al. 2008; Fuentes and Holman 2008; Fraser and Kavelaars 2009; Fuentes et al. 2009).  As the scattered disk and Sedna populations are not close to completion on the large end ($H<4.5$ mags), including such objects (Eris, 2007 OR10 and Sedna), as done here, likely results in a shallower measured slope.  The size distributions of individual dynamical classes are likely more informative (see section 4.2.2).  There are no obvious discontinuities at the large end of the KBO size distribution when including all dynamical classes of TNOs (Figures~\\ref{fig:KBOcumH} and~\\ref{fig:KBOdiffH}).   \\subsubsection{Dwarf Planets}  A dwarf planet is defined by the International Astronomical Union (IAU) as an object that is in hydrostatic equilibrium and has not cleared the neighborhood around its heliocentric orbit of other similarly sized objects.  Though the dwarf planet definition is imprecise, it is clear that Ceres in the main asteroid belt as well as Pluto and Eris in the outer solar system are bonafide dwarf planets. Makemake and Haumea are also likely dwarf planets as are the next largest bodies in the outer solar system such as Sedna, 2007 OR10, Orcus and Quaoar.  Though the lower size limit of an object in hydrostatic equilibrium is not well defined, Lineweaver and Norman (2010) suggest it could be as small as 200 km in radius for an icy body in the outer solar system.  This would put tens more objects in the outer solar system into the dwarf planet category, including three objects discovered in this survey (Table 1: 2010 EK139, 2010 KZ39 and 2010 FX86).  The actual sizes and shapes of these bodies are not well known to date and will depend heavily on their albedos and compositions.  Further detailed observations are required to determine the true sizes and shapes of the new discoveries.  With most of the biggest Kuiper Belt objects likely known, it is interesting to compare where the largest ($H\\leq 4.5$ mags) objects reside dynamically in the Kuiper Belt (Figure~\\ref{fig:kboea2011}). At least one of the largest objects can be found in most of the TNO dynamical populations (Tables 3 and 4).  The scattered disk population (Gomes et al. 2008) has Eris and 2007 OR10, while Sedna is in its own dynamical class (Morbidelli and Levison 2004; Gladman and Chan 2006) that resides significantly beyond the Kuiper Belt edge (Trujillo and Brown 2001; Allen et al. 2001).  The high inclination classical Kuiper belt (Gomes 2003) has several large objects including Makemake, Haumea, Varuna and (278361) 2007 JJ43.  Even the low inclination classical Kuiper belt population, generally known for its smaller sized objects (Levison and Stern 2001), appears to have Quaoar. Further confirming Quaoar's status as a low inclination Kuiper belt object is Quaoar's ultra-red surface (Jewitt and Luu 2004), which is a characteristic generally associated with the low inclination classical Kuiper Belt (Tegler and Romanishin 2000; Trujillo and Brown 2002; Stern 2002; Doressoundiram et al. 2008; Peixinho et al. 2008).  The actual number of Pluto sized bodies is now known (Table 3). Previous authors have argued that the Kuiper Belt likely lost a substantial amount of its mass through collisional grinding and dynamical interactions with the planets (Kenyon and Luu 1999; Levison et al. 2008; Morbidelli et al. 2008; Stewart and Leinhardt 2009). Observationally, many more objects appear to be required in order to produce the observed angular momentum of the largest KBOs (Jewitt and Sheppard 2002; Rabinowitz et al. 2006) and binaries (Noll et al. 2008).  Detailed simulations show that Kuiper Belt formation is possible with only the small number of Pluto sized objects observed (Kenyon and Bromley 2008; Schlichting and Sari 2011).  A significant number of Pluto sized objects likely exist in the populations beyond 100 AU such as the Sedna types and Oort cloud objects, which are currently too faint to be efficiently detected to date.  It is important to determine if the Pluto sized objects formed in the Kuiper Belt as we see it today or if they originated much closer to the Sun and were later transported to their current orbits.  \\subsubsection{TNO Population Ratios}  On the large size end ($H \\lesssim 4.5$ mags), the ratio of the (Plutinos):(Main Kuiper Belt):(Scattered Disk):(Sedna Types) was found to be $(1):(2.6):(7\\pm3):(75\\pm_{-55}^{+115})$, respectively (Table 4).  Thus the Sedna population could be the dominant observed small body population for dwarf sized planets (Figure~\\ref{fig:KBO32}).  The scattered disk population is likely bigger than the main Kuiper Belt (MKB) population by a factor of a few.  The Plutino population is smaller by a factor of a few compared to the main Kuiper Belt.  Both the scattered disk and main Kuiper Belt populations on the large end of the size distribution ($H \\lesssim 4.5$ mags) are consistent with $q = 3.3\\pm0.7$ while the Plutino population appears significantly shallower than this with $q = 2.2\\pm0.5$.  The scattered disk population size determined from the largest objects ($H<4.5$ mags) is consistent with Trujillo et al. (2001) estimated from smaller objects in the scattered disk when using a $q\\sim 3.3$ size distribution.  \\subsubsection{The Main Kuiper Belt}  The main Kuiper Belt ($39 < a < 48$ AU) appears to be divided into three distinct dynamical classes (Figure~\\ref{fig:KBOmkb}).  The high inclination and low inclination (``cold'') classical classes have been suggested for a decade, with the largest objects preferentially in high inclination orbits (Levison and Stern 2001; Brown 2001).  When plotting only the largest few objects ($H<4.5$ mags), there appears to also be both low eccentricity and higher eccentricity classes (Figure~\\ref{fig:kboea2011}).  All three of the low inclination objects ($i<10$ degs) with $H<4.5$ mags have low eccentricities ($e<0.05$).  The high inclination objects with $H<4.5$ mags in the main Kuiper Belt appear to have either low eccentricities ($0.03 < e < 0.07$; 6 observed) or significantly higher eccentricities ($0.13 < e < 0.16$; 8 observed).  Only one of the twenty main Kuiper Belt objects with $H<4.5$ mags has an eccentricity between these two ranges (Salacia (120347) 2004 SB60 which has $e=0.10$) while two others have slightly higher eccentricities (Haumea with $e=0.20$ and (230965) 2004 XA192 with $e=0.25$).  The Hartigan and Hartigan (1985) dip test for bimodality shows a strong bimodality in eccentricity when including all main Kuiper Belt objects with $H<4.5$ mags except for the interesting binary object Salacia (these nineteen objects give a dip statistic of 0.145 which corresponds to a confidence of 0.997 for bimodality, a 3 sigma result).  Including Salacia in the dip test gives a less significant result of only 0.990 confidence in bimodality, or slightly less than 3 sigma.  Including smaller main Kuiper Belt objects decreases the bimodality significance even further.  If real, the low versus higher eccentricity populations of highly inclined large objects could have different origins, such as forming in different regions of the solar system or originally from different scattering events during the migration of the planets.  The largest bodies ($H<4.5$ mags) are too few for meaningful statistics, but it appears that the high inclination main Kuiper Belt does not have a significantly shallower power-law distribution than the low inclination population, as was found for the smaller objects of these two populations by Fraser et al. (2010) (Figure~\\ref{fig:KBOmkb}).  There is a possible deficiency of objects in the main Kuiper Belt between $2.5 < H < 3.5$ magnitudes, but this is likely not statistically significant as it is just small number statistics.  \\subsubsection{Resonance Populations}  A surprising result is the absence of large objects in all the main Neptune resonance populations except the 3:2 resonance (see Gladman et al. 2008 and Elliot et al. 2005 for resonance calculations as well as the updated version of Elliot et al. 2005 kept by Marc Buie at www.boulder.swri.edu/buie/kbo/astrom).  The Plutinos or 3:2 resonance objects include some of the largest known KBOs such as Pluto, Orcus and Ixion (Table 3) while the other observed heavily populated resonances such as the 5:3, 7:4, 2:1, and 5:2 have no known large KBOs (Table 4).  Any object in the Neptune resonances brighter than 21st magnitude ($r \\gtrsim 200$ km), would likely have been detected by now (Figure~\\ref{fig:distanceH} and Table 4).  The 5:2 has a few sizable objects with absolute magnitudes of around 3.8 and 5.1 mags, (84522) 2002 TC302 and (26375) 1999 DE9, respectfully.  The largest 2:1 object appears to be (119979) 2002 WC19 with an absolute magnitude of 5.1 mags.  None of the other resonances have any objects with absolute magnitudes brighter than 5 mags.  2010 EK139, discovered in this survey, appears to be one of the only known objects in the very distant 7:2 resonance (based on orbit calculations from Marc Buie's website at www.boulder.swri.edu/buie/kbo/astrom that has up to date information first published in Elliot et al. 2005).  The relative populations of the various Neptune resonances are currently not well constrained since observational biases make discoveries easier in the closer 3:2 resonance (Jewitt et al. 1998; Trujillo et al. 2001).  There is also a strong longitude and latitude dependence on discovery of resonance populations (Chiang and Jordan 2002; Chiang et al. 2003).  Previous observational works have suggested that the 2:1 resonance appears to have less objects than the 3:2 resonance (Jewitt et al. 1998; Chiang and Jordan 2002).  Numerical simulations of resonance sweeping (Hahn and Malhotra 2005) have shown that the main Neptune resonance populations relative to the 3:2 resonance population may be 2:1 (x2), 7:4 (x0.8), 5:3 (x0.6), 5:2 (x0.5).  These simulations suggest that the outer resonances should have a factor of four more objects than the 3:2 resonance.  Thus, if 3 very large objects such as Pluto, Orcus and Ixion were found in the 3:2 resonance population, based on Poisson statistics, one would expect $12\\pm3.5$ objects of similar size in the other resonances. This is not the case, and such a scenario can be rejected with 3.5 sigma confidence, so either the outer resonances are significantly less populated than the 3:2 resonance or the outer resonance bodies have a different size distribution than the 3:2 resonance.  It is likely that the 3:2 resonance is populated by objects that formed significantly closer to the Sun than the outer resonances.  Objects forming closer to the Sun would likely accrete more material in a shorter amount of time allowing them to become larger before they were captured in the Neptune resonances.  With the Kuiper Belt nearly complete to 21st magnitude, it is unlikely that a planet larger than Mercury within a few 100 AU currently exists.  Lykawka and Mukai (2008) suggested such a planet could have disrupted the outer resonance populations.  It is still possible that a close stellar encounter or now defunct outer planet could have disrupted or depleted the outer resonances early in the solar system's history."
    },
    "1107/1107.3188.txt": {
        "abstract": "We describe a method for  estimating physical conditions in the broad line region (BLR) for a significant subsample of Seyfert-1 nuclei and quasars. Several diagnostic ratios  based on  intermediate (\\aliii, \\siiii)  and   high (\\civ, \\siiv) ionization  lines in the UV spectra of quasars are used to constrain density, ionization and metallicity of the emitting gas. We apply the method to two extreme Population A quasars -- the prototypical NLSy1 I Zw 1 and higher $z$\\ source SDSS J120144.36+011611.6. Under assumptions of spherical symmetry and pure photoionization we infer BLR physical conditions: low ionization (ionization parameter $< 10^{-2}$),  high density (10$^{12} - 10^{13}$ \\cm3) and  significant metal enrichment. Ionization parameter and density can be derived independently for each source with an uncertainty that is  less than $\\pm 0.3$ dex.  We use the product of density and ionization parameter  to estimate the BLR radius and derive an estimation of the virial black hole mass (M$_{BH}$).  Estimates of M$_{BH}$ based on the ``photoionization'' analysis described in this paper are probably  more accurate than those derived from  the mass -- luminosity correlations widely employed to compute black hole masses for high redshift quasars. ",
        "introduction": "We  lack a simple diagnostic method to estimate physical conditions (density, ionization parameter, metallicity) in the broad line region (BLR) of quasars.  Techniques for estimating electron  density and ionization in nebular astrophysics \\citep{osterbrockferland06} are not straightforwardly applicable to broad lines in quasars. Reasons include: large line widths, line doublets too closely spaced to resolve individual components, and density at least an order  of magnitude higher than the critical density assumed for forbidden transitions modeled in spectra of planetary nebul\\ae\\ and \\hii\\ regions. The ionization parameter $U$, that represents the dimensionless ratio of the number of ionizing photons and the electron density \\ne\\ or, equivalently, the total number density of hydrogen \\nh\\, ionized and neutral\\footnote{ In a fully ionized medium \\ne\\ $\\approx 1.2$ \\nh. We prefer to adopt a definition based on \\nh\\ because it is the one employed in CLOUDY computations.}, can be estimated from the intensity ratio of two strong  resonance lines arising from different  ionic stages of the same element \\citep{davidsonnetzer79}.  It is again not easy to interpret the results in the case of quasars. For example, the ratio  \\ciii/\\civ\\ may not yield a meaningful value if the relatively  low \\ciii\\  critical density  implies that the line is  formed at larger radii than \\civ.  Using lines of different ions widens the choice of diagnostic line ratios although additional  sources  of uncertainty are introduced.  The presence of the strong  \\ciii\\  emission line implies that  electron density \\ne\\ cannot be very high. However,  high density (\\ne $ \\sim 10^{11-13}$\\cm3) is invoked to explain the  rich low ionization spectrum (especially \\feii) observed in  quasars \\citep[e.g., ][]{baldwinetal04}.  Several lines in the UV spectrum of I Zw 1 (prominent \\feii, relatively strong \\aliii, and  C{\\sc iii}\\l1176) point towards high density at least for the low ionization line (LIL) emitting zone  \\citep{baldwinetal96,laoretal97a}. This low ionization BLR (LIL BLR) has very similar properties to the O{\\sc i}  and Ca{\\sc ii} emitting region identified by \\citet{matsuokaetal08}. The region where these LILs  are produced cannot emit much \\ciii\\ if electron density exceeds 10$^{11}$ \\cm3.  But is the \\ciii\\  line really so strong  in most quasars?  BLR conditions are certainly complex and the assumption of a single emitting region cannot explain both LILs and high ionization lines (HILs) \\citep[][and references therein]{marzianietal10,wangetal11}.  In this paper we report  an analysis based on several diagnostic ratios used to constrain density, ionization parameter and metallicity in the BLR of two sources that are representative of the Narrow Line Seyfert 1 (NLSy1) sub-sample of quasars (\\S \\ref{sources}). We use methodological considerations that enable us to deblend and identify the principal lines of the spectra  (\\S \\ref{sec:considerations}). Our sources show weak  \\ciii\\ emission (relative to \\siiii) simplifying interpretation of the emission line spectrum (\\S \\ref{sec:measurements}).  Diagnostic ratios are heuristically  defined (\\S \\ref{ratios})  and interpreted through an array of photo-ionization simulations (\\S \\ref{interp}). We show that they converge toward well-defined values of ionization and density (\\S \\ref{results}). We also consider the influence of  heating sources other than photoionization (\\S \\ref{alt}). Under the assumption of a pure photoionized gas, the present analysis can be used to determine the product of density and ionization parameter enabling us to estimate the distance of the BLR from the central continuum source (\\rb) and the virial black hole mass (\\mbh)\\ (\\S \\ref{mass}). Finally, in \\S \\ref{conc} we give our conclusions.  \\subsection{The Eigenvector 1 Parameter Space} \\label{sources}  \\citet{borosongreen92} (BG92) identified a series of correlations from principal component analysis (PCA) of the correlation matrix of emission line measures for a bright low-redshift quasar sample. The sample contained 87 PG quasars with $z \\sim$ 0.5 (17 of them radio-loud, RL).  The first PCA eigenvector (hereafter E1) identified  correlations involving broad  \\hb, broad FeII and narrow \\oiiiopt\\ emission lines.  In an effort to clarify the meaning of the E1, \\citet{sulenticetal00a} searched for a correlation diagram that  showed maximal discrimination between the various Active Galaxy Nuclei (AGN) subclasses. The best E1 correlation space that we could identify involved measures of: 1) equivalent width of the \\feii\\l4570 blend (defined as the ratio \\rfe = W(\\feii\\l4570)/W(\\hb)) and 2) FWHM(\\hb). These were supplemented with: 3) the soft X-ray photon index, \\Gsoft\\ and the centroid line shift of high ionization \\civ. Figure 7 of \\citet{sulenticetal00a} shows two-dimensional (2D) projections of this four-dimensional E1 (4DE1) space: FWHM(\\hb) vs. \\rfe,  \\Gsoft\\ vs. \\rfe\\ and FWHM(\\hb) vs. \\Gsoft. They supplemented the BG92 RL sample with an additional 18 sources (with comparable signal-to-noise S/N spectra) taken from Marziani et al. (1996) who reported W(\\feii) measures over the  range 4240 -- 5850 \\AA.  The  range 4240 -- 5850 \\AA\\ \\feii\\ flux  was divided by a factor $\\approx$ 3.3 in order to obtain the flux in the range 4434-4684 \\AA\\ that was used  by BG92  (both works relied on  the same \\feii\\ template based on I Zw 1).  \\citet{sulenticetal00a} separated the various subclasses of AGN into two populations: Population A-B separated at FWHM(\\hb)= 4000 \\kms. Physical drivers for the correlation were discussed in  \\citet[e.g.][]{marzianietal03b} with black hole mass and Eddington ratio $L/L_{Edd}$ (where $L_{Edd} = 1.5 \\times 10^{38} (M/M_\\odot)$ is the Eddington Luminosity) identified as the principal drivers of change along the 4DE1 sequence. Black hole mass increases from Pop. A to Pop. B  while Eddington ratio decreases from Pop. A to Pop. B.  The division into two populations is, at the least, useful for highlighting major differences among Type 1 AGN, although spectral differences among objects within the same population are still noticeable, especially for Pop A sources (Figure 2 of \\citealt{sulenticetal02}). This is the reason why they also divided the 4DE1 optical plane into bins of $\\Delta$FWHM(\\hb) = 4000 \\kms\\ and $\\Delta$\\rfe = 0.5. Bins A1, A2, A3, A4 are defined in terms of increasing \\rfe,  while bins B1, B1$^+$, and B1$^{++}$ are defined in terms of increasing FWHM(\\hb) (see Fig. 1 of \\citealt{sulenticetal02}).  Sources belonging to the same spectral type show  similar spectroscopic measures and physical parameters (e.g. line profiles  and UV line ratios). Systematic changes are minimized within each spectral type so that an individual quasar can be taken as representative of all sources within a given spectral bin. The binning adopted in \\citet{sulenticetal02} is valid for low $z$\\ ($<$ 0.7) quasars. At higher $z$ an adjustment must be made since no sources with FWHM(\\hb) $\\textless$ 3500 \\kms\\ are found above redshift $z \\sim$ 3 \\citep{marzianietal09}. ",
        "conclusions": "\\label{conc}  Diagnostic line ratios for estimation of density, ionization and metallicity can be found and exploited for  NLSy1-like sources at high redshift. Accurate diagnostics require high S/N and moderate dispersion spectra but in principle can be applied to very high $z$ ($>$ 6.5) using data from IR spectrometers.  The product  (\\nh$\\cdot U$)  yields the possibility of deriving \\rb\\ and \\mbh\\ for a significant number of sources for comparison with estimates derived from extrapolation of the Kaspi relation. Negrete et al. (2012, in preparation) will present an analysis of the applicability of the photoionization method described here to the general population of quasars  starting from the sources with reverberation mapping determinations of \\rb.  It is important to stress however, that the line deblending that allow us to obtain the line fluxes used to compute \\nh\\ and $U$, is not a trivial task. There is a wide diversity in the quasar spectra and one must consider previous results that allow us to predict and/or expect the presence and shape of various components in the broad lines  (as discussed in \\S\\ref{sec:considerations}). One also has to take into account expectations on certain line ratios (e.g. the extreme Pop. A sources show the lowest \\ciii/\\siiii\\ ratio among all quasars in the E1 sequence, see \\S \\ref{sec:ciii}).  The results of this study are preliminary in many ways. Several issues remain open. 1) Is there a ``universal'' \\nh$\\cdot U$ product that can be used for all AGN? The consistency among our results, the ones reviewed in \\S \\ref{neu} and the \\rb-$L$ relation, suggest that this might be the case, at least to a first approximation and especially for Population A sources.  2) A grid of simulations could be refined considering more intermediate metallicity cases. 3) Application of this method has been carried out without considerations about the geometry and kinematics of the BLR. A more refined treatment should consider likely scenarios and investigate their influence on derived physical parameters. 4) Considerations of geometry and kinematics could lead to a physical model accounting for non-thermal heating and production of \\feii\\ emission in a context more appropriate than that of pure photoionization.  As a final remark, we stress that values of \\mbh\\ derived from photoionization considerations, even assuming an average (\\nh$\\cdot U$) are probably more accurate than the ones derived from the mass-$L$\\ relation. The main reason is that we are using each individual quasar luminosity (and not a correlation with large scatter) and the properties of a region that remains similar to itself over a wide range of luminosity."
    },
    "1107/1107.2145_arXiv.txt": {
        "abstract": "{% Among type IIP supernovae there are a few events that resemble the well-studied supernova 1987A produced by the blue supergiant in the Large Magellanic Cloud. }{% We study a peculiar supernova 2000cb and compare it with the supernova 1987A. }{% We carried out hydrodynamic simulations of the supernova in an extended parameter space to describe its light curve and spectroscopic data. The hydrogen H$\\alpha$ and H$\\beta$ lines are modeled using a time-dependent approach. }{% We constructed the hydrodynamic model by fitting the photometric and spectroscopic observations. We infer a presupernova radius of $35\\pm14 R_{\\sun}$, an ejecta mass of $22.3\\pm1 M_{\\sun}$, an explosion energy of $(4.4\\pm0.3)\\times10^{51}$ erg, and a radioactive $^{56}$Ni mass of $0.083\\pm0.039 M_{\\sun}$. The estimated progenitor mass on the main sequence lies in the range of $24-28 M_{\\sun}$. The early H$\\alpha$ profile on Day 7 is consistent with the density distribution found from hydrodynamic modeling, while the H$\\alpha$ line on Day 40 indicates an extended $^{56}$Ni mixing up to a velocity of 8400 km\\,s$^{-1}$. We emphasize that the dome-like light curves of both supernova 2000cb and supernova 1987A are entirely powered by radioactive decay. This is unlike normal type IIP supernovae, the plateau of which is dominated by the internal energy deposited after the shock wave propagation through the presupernova. We find signatures of the explosion asymmetry in the photospheric and nebular spectra. }{% The explosion energy of supernova 2000cb is higher by a factor of three compared to supernova 1987A, which poses a serious problem for explosion mechanisms of type IIP supernovae. } ",
        "introduction": "\\label{sec:intro} Type II-plateau supernovae (SNe IIP) represent the most numerous subclass of core-collapse SNe. They are characterized by a $\\sim100$ day plateau in the light curve, which is a generic feature of the explosion of a red supergiant (RSG) star (Grassberg et al. \\cite{GIN_71}; Falk \\& Arnett \\cite{FA_77}). This picture is in line with the theory of stellar evolution that predicts that stars in the range of $9-25...30~M_{\\sun}$ end their life as RSGs (Heger et al. \\cite{HFWLH_03}). The RSG nature of type IIP pre-SNe was confirmed by the detection of pre-SNe in archival images (Smartt \\cite{Sma_09}).  SN~1987A in the Large Magellanic Cloud (LMC), which was identified with the explosion of a blue supergiant (BSG), became a serious challenge for theoreticians. There is still no clear answer to why the pre-SN was a small-radius star. Two major explanations were proposed: (i) the mixing between the He-core and the hydrogen envelope mediated by a fast rotation in combination with the low metallicity (Saio et al. \\cite{SNK_88}; Weiss et al. \\cite{WHT_88}; Woosley et al. \\cite{WPE_88}); (ii) the loss of an extended hydrogen envelope in a close binary system (Hillebrandt \\& Meyer \\cite{HM_89}; Podsiadlowski \\& Joss \\cite{PJ_89}). The unresolved issue of the pre-SN origin in the case of SN~1987A is a strong stimulus for studying objects that resemble one another. Pastorello et al. (\\cite{PBB_05}) studied a peculiar SN~1998A, another example of the SN originating in the explosion of a BSG. This SN is very similar photometrically and spectroscopically to SN~1987A, although its explosion energy probably was somewhat higher than that of SN~1987A (Pastorello et al. \\cite{PBB_05}).  Recently, Kleiser et al. (\\cite{KPK_11}) have analyzed photometric and spectroscopic observations of SN~2000cb and conclude that the SN was related with the explosion of a BSG. In that sense SN~2000cb is another counterpart of SN~1987A. Nevertheless, authors have emphasized differences, particularly, in the light curve shape and in expansion velocities, which were substantially higher in SN~2000cb. Using scaling relations and SN~1987A as a template, Kleiser et al. (\\cite{KPK_11}) have derived the ejecta mass of $\\sim16.5~M_{\\sun}$ and the explosion energy of $\\sim4\\times10^{51}$ erg for SN~2000cb. On the other hand, hydrodynamic modeling has led Kleiser et al. (\\cite{KPK_11}) to the modest estimate of $\\sim1.7\\times10^{51}$ erg, comparable to that of SN~1987A. It should be noted, however, that authors did not attain a reasonable fit of the bolometric light curve. The question of reliable SN~2000cb parameters therefore remains open.  Motivated by the importance of the SN~1987A-like events and rather complete observational light curve of SN~2000cb (Kleiser et al. \\cite{KPK_11}), we revisit the analysis of this SN by employing a standard hydrodynamic model (Utrobin \\cite{Utr_04}, \\cite{Utr_07}). A brief description of the hydrodynamic model is presented in Sect.~\\ref{sec:hydmod}. Because the SN parameters cannot be recovered from the photometric data alone, one also needs to use the line profiles in which the mass-velocity distribution is imprinted. By this approach we obtain the basic parameters of the optimal model given in Sect.~\\ref{sec:result-param}, while in Sect.~\\ref{sec:result-asym} we address signatures of the explosion asymmetry in the photospheric and nebular spectra. In Sect.~\\ref{sec:disc} we discuss results and their implications for the explosion mechanism problem, and in Sect.~\\ref{sec:concl} we summarize the conclusions.  Following Kleiser et al. (\\cite{KPK_11}) we adopt a distance to the spiral galaxy IC 1158 of 30 Mpc, a reddening $E(B-V)=0.114$ mag, an explosion date of April 21.5 UT (JD 2\\,451\\,656), and a recession velocity to the host galaxy of 1918.67 km\\,s$^{-1}$. ",
        "conclusions": "\\label{sec:concl} Our goal was to derive the parameters of the peculiar type IIP SN~2000cb from the available observational data by employing hydrodynamical modeling and calculating hydrogen line profiles. We conclude that SN~2000cb was the high-energy explosion of a massive BSG, which was more energetic than SN~1987A by a factor of three. At the same time, these events are close counterparts by their pre-SN structure and by the common physics of their dome-like light curves. SN~2000cb can thus be considered as a high-energy version of SN~1987A.  The progenitor mass, the explosion energy, and the radioactive $^{56}$Ni mass of SN~2000cb are consistent with the correlations revealed by the well-observed SNe~IIP studied hydrodynamically. This combined with the clear evidence of one-sided $^{56}$Ni ejecta, which is rather ubiquitous among SNe~IIP, favors a common explosion mechanism for all these SNe. We stress that the high explosion energy of SN~2000cb is far beyond the possibilities of the present-day versions of the neutrino-driven or magneto-rotational mechanism."
    },
    "1107/1107.0971_arXiv.txt": {
        "abstract": "{Low-mass star-forming cores differ from their surrounding molecular cloud in turbulence, shape, and density structure. } {We aim to understand how dense cores form out of the less dense cloud material by studying the connection between these two regimes. } {We observed the L1517 dark cloud in C$^{18}$O(1--0), N$_2$H$^+$(J=1--0), and SO(J$_{\\mathrm{N}}$=3$_2$--2$_1$) with the FCRAO 14m telescope, and in the 1.2mm dust continuum with the IRAM 30m telescope. } {Most of the gas in the cloud lies in four filaments that have typical lengths of 0.5~pc. Five starless cores are embedded in these filaments and have chemical compositions indicative of different evolutionary stages. The filaments have radial profiles of C$^{18}$O(1--0) emission with a central flattened region and a power-law tail, and can be fitted approximately as isothermal, pressure-supported cylinders. The filaments, in addition, are extremely quiescent. They have subsonic internal motions and are coherent in velocity over their whole length. The large-scale motions in the filaments can be used to predict the velocity inside the cores, indicating that core formation has not decoupled the dense gas kinematically from its parental material. In two filaments, these large-scale motions consist of oscillations in the velocity centroid, and a simple kinematic model suggests that they may be related to core-forming flows. } {Core formation in L1517 seems to have occurred in two steps. First, the subsonic, velocity-coherent filaments have condensed out of the more turbulent ambient cloud. Then, the cores fragmented quasi-statically and inherited the kinematics of the filaments. Turbulence dissipation has therefore occurred mostly on scales on the order of 0.5~pc or larger, and seems to have played a small role in the formation of the individual cores. } ",
        "introduction": "Dense cores in nearby dark clouds are the birth sites of solar-mass stars, and they represent the simplest environments where stars can be formed. The study of the internal structure of these cores offers a unique opportunity to determine the initial conditions of low-mass star formation, and for this reason an intense observational effort is under way to characterize cores in detail (see recent reviews by \\citealt{DIF07,WAR07,BER07}).  Compared to their turbulent parent clouds dense cores are very quiescent. Molecular line observations in tracers like ammonia reveal that the gas in a dense core has subsonic internal motions, indicating that thermal pressure contributes more than turbulence in providing support against gravity \\citep{MYE83_2}. In addition, the gas motions inside a core are ``coherent,'' in the sense that the observed linewidth does not depend on radius \\citep{GOO98}, and thus deviates from the linewidth-size relation characteristic of the large-scale gas in a cloud \\citep{LAR81}. Also in contrast to the irregularly shaped cloud gas, cores often display a regular, close-to-spherical geometry \\citep{MYE91}. A number of cores even present centrally concentrated density profiles close to those predicted by hydrostatic equilibrium models \\citep{ALV01}, again suggesting that the core gas conditions differ significantly from those of the cloud gas on larger scales.  How quiescent, centrally-concentrated cores form out of the more turbulent, less-dense cloud material is still a matter of debate. A number of core-formation mechanisms have been proposed over the years, ranging from the quasi-static loss of magnetic support via ambipolar diffusion \\citep{SHU87,MOU99} to the rapid dissipation of turbulence due to supersonic shocks \\citep{PAD01,KLE05,VAZ05}. Observationally, it is unclear which of these proposed scenarios can fit the variety of existing constraints. Ambipolar diffusion models, for example, seem to predict stronger magnetic fields than observed \\citep{CRU10}, and they usually predict contraction times that significantly exceed those inferred from observed core lifetimes and chemical clocks \\citep{LEE99b,TAF02}. Models of core formation by shocks, on the other hand, predict linewidths and velocity displacements between tracers that are larger than commonly observed in cores \\citep{WAL04,KIR_H07}.  Observational progress in constraining the mechanism of core formation requires probing the connection between the dense cores and the less-dense material that surrounds them. This less-dense material likely represents the gas out of which the cores have condensed, so by comparing its geometry, density, and kinematics with those of the core gas, we could infer the physical changes involved in the formation of a dense core. Unfortunately, studying the cloud-to-core transition requires combining observations of tracers sensitive to different density regimes, and until recently, such multi-tracer analysis has been hindered by a number of inconsistencies between the emission from different molecules \\citep{ZHO89,LEM96}.  Work carried out over the past decade has revealed that most inconsistencies between tracers result from chemical changes that occur in the gas at dense-core densities and, in particular, from the selective freeze-out of molecules onto cold dust grains \\citep{KUI96,CAS99,BER02,TAF02,AIK05}. As a result of this work, a reasonably consistent picture of the chemical behavior of the different gas tracers has emerged. According to this picture, most carbon-bearing molecules (including CO and CS) disappear from the gas phase at densities close to a few $10^4$~cm$^{-3}$, while nitrogen-bearing molecules like NH$_3$ and N$_2$H$^+$ survive undepleted up to densities that are at least one order of magnitude higher. In addition, several species like SO and C$_2$S present significant abundance enhancements at early evolutionary times, but end up freezing out onto the grains at similar densities to the C-bearing species (see \\citealt{BER07} for a review of core chemistry).  This new understanding of molecular chemistry at intermediate gas densities now makes it possible to combine observations of different tracers and to finally explore the transition from cloud gas to core material self-consistently. An excellent region for investigating this transition is the L1517 dark cloud in the Taurus-Auriga molecular complex (see \\citealt{KEN08} for a Taurus overview). L1517 appears in optical images as a region of enhanced obscuration associated with the reflection nebulosity from the young stars AB Aur and SU Aur \\citep{LYN62,STR76,SCH79}. The CO observations of the cloud have shown that the extended gas consists of several filamentary components that extend to the northwest of the nebulosity and occupy a region of about $20' \\times 10'$ coincident with the optical obscuration \\citep{HEY87}. Embedded in these components lies a collection of several dense cores that are easily distinguished in the optical images thanks to the contrast provided by the bright nebulosity from AB Aur and SU Aur, which lie at least 0.3~pc in projection from the cores \\citep{SCH79}. Radio observations of these cores reveal the physical properties typical of the Taurus-Auriga core population \\citep{BEN89}, and a prominent member of the group, L1517B, has been the subject of detailed study owing to its regular shape and clear pattern of molecular freeze out \\citep{TAF04,TAF06}.  A notable feature of the L1517 cores is that they all appear to be starless \\citep{STR76,BEI86,KIR07}. This lack of embedded protostars is probably responsible for the L1517 cloud presenting some of the narrowest line profiles in CO and dense gas tracers seen towards the Taurus-Auriga region, and indicates that, although the bright PMS star AB Aur ($\\sim 50$~L$_\\odot$, \\citealt{VAN98}) is physically related to the cloud \\citep{NAC79,DUV86}, the bulk of the L1517 material and its embedded cores remain unperturbed by the energetic stellar output \\citep{LADD91}. This combination of quiescent state, compact size, and multiple starless cores makes the L1517 cloud an ideal laboratory for studying the process of core formation. In this paper, we present a study of the physical conditions and kinematics in both the dense cores and the less-dense surrounding material using a variety of molecules known to trace different density regimes and chemical evolutionary stages of the cloud gas. As will be seen, our analysis suggests that the cores in L1517 have formed by the gravitational contraction of subsonic, velocity coherent gas in 0.5~pc-long filaments. ",
        "conclusions": "We observed the L1517 dark cloud in C$^{18}$O(1--0), N$_2$H$^+$(1--0), SO(J$_{\\mathrm{N}}$=3$_2$--2$_1$) with the FCRAO telescope, and in the 1.2mm dust continuum with the IRAM 30m telescope. From the analysis of these data, we came to the following main conclusions.  1. The gas in the L1517 cloud is structured in four filaments with typical sizes of about 0.5~pc.  2. The radial profile of C$^{18}$O(1--0) emission in each filament consists of a central flattened region and a power-law tail. An isothermal cylinder model provides an approximate fit to the emission, although better agreement is obtained with a softened power law profile.  3. Five starless cores are embedded in the filaments. Their chemical composition indicates that they are at different evolutionary stages, although their kinematic properties are very similar.  4. The velocity field of the gas in the L1517 cloud has been characterized by fitting Gaussians to the observed spectra. Most positions require one Gaussian, although a few regions require two, very likely due to the overlap between different components.  5. The filaments are extremely quiescent. Their nonthermal linewidth is subsonic and changes very little over the length of the filaments. The velocity centroids also change subsonically over the filament length. These characteristics indicate that the gas in the filaments is velocity coherent on scales of about 0.5~pc.  6. Although quiescent, the filaments have large-scale patterns of (subsonic) velocity. The gas in the dense cores follows the velocity pattern of the less dense gas  closely, indicating that core formation has not decoupled the dense gas kinematically from its surrounding filament material.  7. In two filaments, the large-scale velocity patterns seem to consist of  oscillations. A simple kinematic model shows that, at least in one filament, the oscillatory pattern is consistent with core-forming motions along the axis of the filament.  8. Core formation in L1517 seems to have occurred in two steps. First, the subsonic, velocity-coherent filaments have condensed out of the more turbulent ambient cloud. Then, the cores have fragmented quasi-statically and inherited the kinematics of their parental filament. Turbulence dissipation has therefore occurred on scales of 0.5~pc or larger, and seems to have played little role in the formation of the individual cores."
    },
    "1107/1107.2373_arXiv.txt": {
        "abstract": "Given some assumptions it is possible to derive the most general post-general relativistic theory of gravity for the distant field of a point mass. The force law derived from this theory contains a Rindler term in addition to well-known contributions, a Schwarzschild mass and a cosmological constant. The same force law recently was confronted with solar system precision data. The Rindler force, if present in Nature, has intriguing consequences for gravity at large distances. In particular, the Rindler force is capable of explaining about 10\\% of the Pioneer anomaly and simultaneously ameliorates the shape of galactic rotation curves. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.3280_arXiv.txt": {
        "abstract": "{} {Violent gravitational interactions can change the morphologies of galaxies and, by means of merging, transform them into elliptical galaxies. We aim to investigate how they affect the evolution of galactic magnetic fields. } {We selected 16 systems of interacting galaxies with available VLA archive radio data at 4.86 and 1.4\\,GHz and compared their radio emission and estimated magnetic field strengths with their star-forming activity, far-infrared emission, and the stage of tidal interaction. } { The estimated mean of total magnetic field strength for our sample of interacting galaxies is $14\\pm 5\\,\\mu$G, which is larger than for the non-interacting objects. The field regularity (of $0.27\\pm0.09$) is lower than in typical spirals and indicates enhanced production of random magnetic fields in the interacting objects. We find a general evolution of magnetic fields: for weak interactions the strength of magnetic field is almost constant ($10-15\\,\\mu$G) as interaction advances, then it increases up to $2\\times$, peaks at the nuclear coalescence ($25\\,\\mu$G), and decreases again, down to $5-6\\,\\mu$G, for the post-merger remnants. The main production of magnetic fields in colliding galaxies thus terminates somewhere close to the nuclear coalescence, after which magnetic field diffuses. The magnetic field strength for whole galaxies is weakly affected by the star formation rate (SFR), while the dependence is higher for galactic centres. We show that the morphological distortions visible in the radio total and polarized emission do not depend statistically on the global or local SFRs, while they do increase (especially in the polarization) with the advance of interaction. The constructed radio-far-infrared relations for interacting and non-interacting galaxies display a similar balance between the generation of cosmic rays, magnetic fields, and the production of the thermal energy and dust radiation. } { The regular magnetic fields are much more sensitive to morphological distortions induced by tidal interactions than are the random fields. As a result the polarized emission could be yet another indicator of an ongoing merging process. The found evolution of magnetic field with advancing interaction would definitely imply a stronger effect of magnetic fields on the galaxy surroundings in the earlier cosmological epochs. The process of strong gravitational interactions can efficiently magnetize the merger's surroundings, having a similar magnetizing effect on intergalactic medium as supernova explosions or galactic winds. If interacting galaxies generate some ultra-high energy cosmic rays (UHECRs), the disk or magnetized outflows can deflect them (up to $23\\degr$), and make an association of the observed UHECRs with the sites of their origin very uncertain. } ",
        "introduction": "\\label{s:intro}  Both observational evidence and theoretical considerations show that galaxy interactions are the most important processes during the formation and evolution of galaxies (Struck \\cite{struck99}). This is the case both locally in the nearby universe and, in particular, cosmologically at high redshifts, where galaxy encounters are expected to be more frequent. During such encounters, direct collisions between stars are highly unlikely, while collisions between the interstellar medium (ISM) of the two disks are inevitable and may trigger star formation. The interaction-induced starbursts, nuclear activity, feedback, gas infalls, and outflows, all play critical roles in shaping ISM and the surrounding intergalactic medium (IGM) (Smith et al. \\cite{smith10}).  Some low-speed collisions lead to merging of galaxies. Gas-rich mergers could reach the phase of a luminous infrared galaxy (LIRG), or even of an ultraluminous infrared galaxy (ULIRG) (Sanders \\& Mirabel \\cite{sanders96}). Violent galaxy collisions are also suggested as a source of large-scale shocks and one of possible sources of ultra high energy cosmic rays (UHECR; e.g. Bhattacharjee \\& Sigl \\cite{bhattacharjee00}). However, not all galaxy interactions will result in a merger, and it is equally important to study different types of galaxy interaction to understand their role in galaxy evolution.  There were several systematic attempts to study the effects of collisions on structures of galaxies and the properties of different ISM phases. An early work of Toomre \\& Toomre (\\cite{toomre72}) showed that numerical simulations quite readily reproduce the observed morphology of a wide range of interacting galaxies and explain the formation of bars, tidal arms, and bridges. Toomre (\\cite{toomre77}) compiled a sample of 11 ongoing merger systems at various stages of interaction, which became known as the {\\em Toomre sequence}. The sequence itself and numerical simulations of its members explain how elliptical galaxies can be formed by the coalescence of spiral galaxies. The Toomre sequence was studied over the years in great detail in various wavelength bands in order to understand the evolving properties of interacting systems. It was shown that in moving along the sequence from early to late-stage mergers, a larger fraction of atomic gas appears outside the optical bodies (Hibbard \\& Gorkom \\cite{hibbard96}). In the final stages of nuclear coalescence, there is little \\ion{H}{i} within the remnant bodies, while the tidal material flows in and is able to trigger a nuclear starburst. In accordance with that, Hummel et al. (\\cite{hummel90}) found that the central sources of interacting galaxies are by a factor of five times brighter in the radio wavelengths than those in the isolated spirals.  The studies of broadband (V, I) and narrowband (H$\\alpha$) images obtained with the Hubble Space Telescope (HST) by Laine et al. (\\cite{laine03}) show surprisingly little evidence of any distinct trends in nuclear properties along the Toomre merger sequence. There is just a rise in the nuclear luminosity density in the most-evolved members of the sequence. However, the near-infrared (NIR) observations with the HST (Rossa et al. \\cite{rossa07}), which are less affected by dust obscuration show a distinct trend for the nuclei in the sequence to become more luminous as the merger process advances. In a different way than the non-interacting objects, the Toomre-sequence galaxies also reveal newly formed stars that are more concentrated toward their centres. Rossa et al. also argue that if left to evolve and fade for several Gyrs, the merger remnants would have the $K$-band luminosity profiles of normal elliptical galaxies.  More complicated evolution during the merging process was identified by Brassington et al. (\\cite{brassington07}) in the X-rays observed with {\\it Chandra}. They studied a sample of nine interacting and post-merger objects. The original Toomre sequence was extended in this work to include very old (up to a 3\\,Gyr) merger-remnants. First, they find that unlike H$\\alpha$ and optical emission, the X-ray luminosity peaks about 300\\,Myr before the nuclear coalescence. The subsequent drop in X-ray emission is likely due to large-scale diffuse outflows coming out of the galactic disks and reducing the hot gas density. Second, about 1 Gyr after the coalescence the remnants are X-ray-fainter than typical elliptical galaxies, which are known to be generally strong X-ray emitters. Thirdly, the dynamically older mergers seem to rebuild the hot galactic halo and rise again the X-ray luminosity. This can be explained by an outflowing wind driven into the halo by a declining population of Type Ia supernova (Brassington et al. \\cite{brassington07}).  Although magnetic field constitute an important ingredient of the ISM, little is known about how galaxy interactions affect its strength and evolution, and how they respond to the evolution of other ISM phases presented above . The most detailed analysis of magnetic field in disks of interacting galaxies was done by Chy\\.zy \\& Beck (\\cite{chyzy04}) for the Antennae galaxies. They show that the mean total magnetic field strength (about $20\\,\\mu$G) is significantly greater than in isolated grand-design spirals and may be due to the interaction-induced star formation. The field regularity (the ratio of regular to random field components) appeared to be lower than in undistorted galaxies, which can indicate enhanced turbulent motions in the merger system. A strong regular component (of $10\\,\\mu$G) was found to trace gas-shearing motions along the tidal tail. When and how the enhanced magnetic field arose and when it would disperse in this system are still not known.  The evolution of magnetic field in merging systems like the Antennae may look different from the evolution of the X-ray component revealed by Brassington et al. (\\cite{brassington07}). In general, the early-type galaxies are typically radio weak, and strong radio emission is only observed in ellipticals with an AGN activity. In such active galaxies, generation of strong magnetic fields, although possible in accretion disks, is not connected with any large-scale or small-scale galactic dynamo processes, which operate in typical spiral or irregular galaxies (Beck et al. \\cite{beck96}, Brandenburg \\& Subramanian \\cite{brandenburg05}, Chy\\.zy et al. \\cite{chyzy11}). The first process requires a strong differential rotation with uniform sign of turbulence twisting, the second depends on the ISM turbulence, which likely declines in early-type galaxies owing to a weaker star formation.  The goal of this paper is to investigate the evolution of magnetic field during galaxy interactions. As magnetic fields are best studied from the synchrotron emission, we restrict our considerations to objects with available radio data. The polarized properties of radio emission also provides a unique diagnostic tool for recognizing gas flows and their distortions. In Sect.~\\ref{s:observations} we describe the selection criteria of the constructed sample of interacting galaxies and present the used (VLA\\footnote{The VLA of the NRAO is operated by Associated Universities, Inc., under cooperative agreement with the NSF}) radio data and the details of their reduction. The next section provides a brief description of each interacting system and its properties in radio total and polarized emission. We present estimations of equipartition magnetic field strength for all the galaxies and compare them with those for isolated spirals (Sect. \\ref{s:magnetic}).  The statistical analysis of all the galaxies begins from investigating the evolution of magnetic fields as the interaction and merging process advance (\\ref{s:evolution}). In Sect. \\ref{s:sfr} we examine the relation of magnetic fields to the star formation rate (SFR). An analysis of observed morphological asymmetries in total and polarized radio emission is provided in Sect. \\ref{s:asymmetries}. The radio-far-infrared (FIR) diagram for interacting galaxies is presented in Sect. \\ref{s:radiofir} and compared to the literature results. We also discuss the possible magnetization of IGM by magnetic fields of merger origin and the propagation effects of UHECRs in the vicinity of colliding galaxies (Sect \\ref{s:impact}). We summarize our findings in Sect.~\\ref{s:summary}.  \\begin{table*} \\caption{Interacting systems and their members.} \\begin{center} \\begin{tabular}{lccccccc} \\hline\\hline Name                 & System name         &     Type           & Dist. & Inkl. & Pos. angle & \\ion{H}{i} extent & Interaction \\\\ & or other name       &                    & [Mpc] & [deg.]&  [deg]     & [kpc]             & Stage       \\\\ \\hline \\object{NGC\\,876}$^S$         & NGC\\,876/877        & SAc: sp            & 50.8  & 77.9  & 27.3       &  N/A              & $-1$        \\\\ \\object{NGC\\,877}$^S$         & NGC\\,876/877        & SAB(rs)bc LIRG     & 50.8  & 35.6  & 138.0      &                   & $-1$        \\\\ \\object{NGC\\,4254}            & The Virgo Cluster   & SA(s)c,LINER,HII   & 17    & 32.0  & 60.0       & 40                & $-1$        \\\\ \\object{NGC\\,2207}$^S$        & NGC\\,2207/IC\\,2163  & SAB(rs)bc pec      & 35.0  & 58.2  & 115.9      &  56               & $-1$        \\\\ \\object{IC\\,2163}$^S$         & NGC\\,2207/IC\\,2163  & SB(rs)c pec        & 35.0  & 78.5  & 102.6      &                   & $-1$        \\\\ \\object{NGC\\,5426}            & Arp\\,271, NGC\\,5426/5427 & SA(s)c pec    & 26.7  & 69.7  & 0.5        & 43                & $-1$        \\\\ \\object{NGC\\,5427}$^S$        & NGC\\,5426/5427      & SA(s)c pec,Sy2,HII & 26.7  & 25.5  & 135.0      &                   & $-1$        \\\\ \\object{NGC\\,6907}$^S$        & NGC\\,6907/6908      & SB(s)bc            & 44.5  & 37.5  & 66.7       & 78                & $0$        \\\\ \\object{NGC\\,6908}            & NGC\\,6907/6908      & S                  & 44.5  &  N/A    & N/A      &                   & $0$        \\\\ \\object{UGC\\,12914}$^S$       & Taffy galaxy        & SAB(rs)c,Sbrst,Sy2 & 59.6  & 54.1  & 159.6      & 52                & $1$        \\\\ \\object{UGC\\,12915}           & Taffy galaxy        & Sdm                & 59.6  & 72.9  & 135.4      &                   & $1$        \\\\ \\object{UGC\\,813}             & Taffy2 galaxy       & Sb                 & 67    & 72.0  & 110.3      & 73                & $1$        \\\\ \\object{UGC\\,816}             & Taffy2 galaxy       & Sc                 & 67    & 62.0  & 170.0      &                   & $1$        \\\\ \\object{NGC\\,660}             &                     & SB(s)a pec,HII     & 12.3  & 78.8  & 11.9       & 47                & $1$        \\\\ \\object{NGC\\,4038}$^{STX}$    & Arp\\,244,The Antennae & SB(s)m pec       & 26.8  & 51.9  & 133.2      & 133               & $1$        \\\\ \\object{NGC\\,4039}$^{STX}$    & The antennae        & SA(s)m,pec,LINER   & 26.8  & 71.2  & 132.0      &                   & $1$        \\\\ \\object{NGC\\,6621}$^T$        & Arp\\,81             & Sb pec,LIRG        & 86.4  & 70.8  & 142.5      & 63                & $5$        \\\\ \\object{NGC\\,6622}$^T$        & Arp\\,81             & Sa                 & 86.4  & 27.3  & 117.0      &                   & $5$        \\\\ \\object{NGC\\,520}$^{ST}$      & Arp\\,157            & pec                & 30.2  & 77.4  & 130.0      & 103               & $7$        \\\\ \\object{NGC\\,3256}$^{TX}$     &                     & Sb(s) pec,LIRG     & 56    & 48.7  & 83.2       & 127               & $9$        \\\\ \\object{Arp\\,220}$^X$         &                     & S,Sy,ULIRG         & 76    & 57.0  & 96.5       & 133               & $10$        \\\\ \\object{NGC\\,7252}$^{TX}$     & Arp\\,226,The Atoms of Peace  & (R)SA(r)  & 63    & 25.1  & 127.0      & 214               & $11$        \\\\ \\object{Arp\\,222}$^X$         &                     & SAB(s)a pec        & 23    & 43.4  & 66.5       & N/A               & $12$        \\\\ \\object{NGC\\,1700}$^X$        &                     & E4                 & 54    & 90.0  & 87.0       & N/A               & $13$        \\\\ \\hline \\end{tabular} \\end{center} {\\bf Notes.} $^{(T)}$ - a member of the Toomre sequence (\\cite{toomre77}); $^{(X)}$ - a member of the X-ray sample of Brassington et al. (\\cite{brassington07}); $^{(S)}$ - a member of our compiled sample of angularly-large interacting galaxies \\label{t:sample} \\end{table*} ",
        "conclusions": "\\label{s:summary}  We studied the influence of gravitational interactions on the distribution of radio emission and the properties of magnetic fields in galaxies. We perused archival radio data from the VLA at 4.86\\,GHz and 1.4\\,GHz for a sample of 24 galaxies in 16 interacting systems. Different stages of interaction were approximately described by the interaction stage ($IS$) parameter (Sect. \\ref{s:discussion}). The main conclusions are as follows.  \\begin{itemize}  \\item The estimated total magnetic field strengths are from 5\\,$\\mu$G (ARP\\,222) to 25-27\\,$\\mu$G (NGC\\,3256, ARP\\,220), with a mean of $14\\pm 5\\,\\mu$G, greater than for bright non-interacting galaxies. The mean field regularity of $0.27\\pm0.09$ is less than for typical spirals, which indicates enhanced production of random magnetic fields in the interacting objects.  \\item Detailed analysis of NGC\\,2207/IC\\,2163 reveals a region of enhanced radio total and polarized emission in the eastern part of NGC\\,2207, which can suggest a supply of CR electrons originating in IC\\,2163. Our studies of polarized properties of this system contradict the earlier hypothesis that the enhanced radio emission in this region is due to compression of magnetic fields (Elmegreen et al. \\cite{elmegreen95a}).  \\item The structure of magnetic field $\\vec{B}$-vectors in the polar-ring galaxy NGC\\,660 resembles an X-shape pattern, typical of late-type edge-on galaxies, which challenges some earlier suggestions that this is a post-merger system.  \\item Even in weakly interacting galaxies, the polarized emission and the orientation of magnetic field $\\vec{B}$-vectors are very sensitive tools for revealing global galactic distortions (Taffy galaxies) or the direction of gas flow in tidally stretched spiral arms and tidal tails (NGC\\,2207, NCG\\,6907, NGC\\,4038).  \\item We built a tentative scenario of magnetic field evolution in interacting galaxies. For weak interactions the strength of magnetic field is almost constant ($10-15\\,\\mu$G), then it increases up to $2\\times$ with advancing of interaction to later stages, and decreases again, down to the limiting value of $5-6\\,\\mu$G, for old post-mergers. We suggest that the main production of magnetic fields in colliding galaxies must terminate close to the nuclear coalescence, after which magnetic field diffuses or is kept at the level of turbulent energy of ISM. The revealed evolution of magnetic fields agrees with the scenario of formation of ellipticals from mergers (Toomre \\& Toomre \\cite{toomre72}),  \\item Magnetic field strength is weakly regulated by the galactic star-forming activity (e.g. SFR), showing a wide spread of estimated strength values (with a slope $\\alpha=0.21\\pm 0.02$ and correlation $\\rho=0.63$). The dependence is stronger in galactic centres and weaker outside. There is also a hint that the regular field is anticorrelated with SFR ($\\rho=-0.55$).  \\item Asymmetry in distribution of polarized emission is stronger ($A_P=0.75$) than in total radio intensity ($A_T=0.64$), which can indicate that the regular component of magnetic fields is much more sensitive to morphological distortions induced by tidal interactions than the random magnetic field. This means that, the asymmetry in the radio polarized emission could be yet another indicator of an ongoing merging process. The asymmetries visible in the radio emission do not statistically depend on the global or local SFRs, while  they increase (especially in the polarization) with advancing interaction.  \\item The constructed radio-FIR relation for interacting and non-interacting galaxies shows that the tidally induced intense star formation, morphological distortions, and processes mutually transforming magnetic field components (e.g. by compression and shearing) still lead to producing magnetic energy and cosmic rays that are balanced by the production of thermal energy and dust radiation. The interacting galaxies constitute the group of statistically strongest radio and FIR emitters. We also argued that the configuration of ordered magnetic fields in the bridge of the Taffy system indicates that the modelling of Lisenfeld \\& Voelk (\\cite{lisenfeld10}), adopted to account for the strong radio emission in the bridge, might not be applicable here.  \\item The process of strong gravitational interactions can efficiently magnetize the merger surroundings (likely up to about 200\\,kpc), exerting a similar impact on IGM to its magnetization by supernova explosions and galactic winds. However, tidal interactions are likely insufficient to fill IGM with magnetic fields with a large volume filling factor of $\\approx60\\%$ as recently suggested by Dolag et al. (\\cite{dolag10}). If interacting galaxies generate some UHECRs by large-scale shocks or AGNs, their propagation through the disk or magnetized outflows (bridges or tidal tails) can deflect them up to $23\\degr$ and make association of observed UHECRs with the sites of their origin less obvious.  \\end{itemize}  Among the interacting systems, those at the final stages of merging as well as the post-merging systems (as elliptical galaxies) are the least studied ones with respect to magnetic field evolution. After the interaction-induced starburst declines, the production of cosmic ray electrons ceases and cannot further enlighten magnetic fields (Sect \\ref{s:evolution}). Such fields may still exist in the mergers, before they disperse because of turbulent diffusivity, but could only be discovered by the Faraday rotation methods applied to the background sources. Thus there is a need for very high-resolution and sensitive radio spectropolarimetric data from interferometers, not yet available. Future EVLA, LOFAR and SKA instrumentation could reveal for the first time such fading or relic magnetic fields, if they are actually present in the post-merging systems."
    },
    "1107/1107.3555_arXiv.txt": {
        "abstract": "To understand how best to use observations of Type Ia supernovae (SNe~Ia) to obtain precise and accurate distances, we investigate the relations between spectra of SNe~Ia and their intrinsic colors.  Using a sample of 1630 optical spectra of 255 SNe, based primarily on data from the CfA Supernova Program, we examine how the velocity evolution and line strengths of \\ion{Si}{2} $\\lambda 6355$ and \\ion{Ca}{2} H\\&K are related to the $B-V$ color at peak brightness.  We find that the maximum-light velocity of \\ion{Si}{2} $\\lambda 6355$ and \\ion{Ca}{2} H\\&K and the maximum-light pseudo-equivalent width of \\ion{Si}{2} $\\lambda 6355$ are correlated with intrinsic color, with intrinsic color having a linear relation with the \\ion{Si}{2} $\\lambda 6355$ measurements.  \\ion{Ca}{2} H\\&K does not have a linear relation with intrinsic color, but lower-velocity SNe tend to be intrinsically bluer.  Combining the spectroscopic measurements does not improve intrinsic color inference.  The intrinsic color scatter is larger for higher-velocity SNe~Ia --- even after removing a linear trend with velocity --- indicating that lower-velocity SNe~Ia are more ``standard crayons.''  Employing information derived from SN~Ia spectra has the potential to improve the measurements of extragalactic distances and the cosmological properties inferred from them. ",
        "introduction": "\\label{s:intro}  Type Ia supernovae (SNe~Ia) are very good distance indicators after making an empirical correction based on their light-curve shape and color \\citep{Phillips93, Riess96}.  Using these relations, large SN~Ia samples have a precision of \\about 8\\% for distance estimates \\citep[e.g.,][]{Hicken09:lc}.  This level of precision is adequate to determine that the expansion of the Universe is accelerating \\citep{Riess98:Lambda, Perlmutter99}, and constrain the equation-of-state parameter of dark energy \\citep{Wood-Vasey07, Riess07, Hicken09:de, Kessler09, Amanullah10, Conley11, Sullivan11}. SNe~Ia are even better distance indicators in the rest-frame near-infrared (NIR) \\citep{Mandel09, Mandel11}; however, low-redshift NIR samples are small and high-redshift NIR samples do not yet exist.  Decades ago, \\citet{Branch87} showed that not all SNe~Ia have the same ejecta velocity, as probed by the blueshift of spectral features. \\citet{Branch88} noted that a sample of SNe~Ia had a very broad distribution of \\ion{Si}{2} $\\lambda 6355$ velocity near maximum brightness and concluded: ``SNe~Ia are not {\\it observationally} homogeneous.  The only way to maintain that they are {\\it physically} homogeneous would be to postulate that they are identical but asymmetrical.''  \\citet{Benetti05} showed that the velocity gradient for the \\ion{Si}{2} $\\lambda 6355$ feature did not correlate with light-curve shape for spectroscopically normal SNe~Ia.  This observation provided further proof that SNe~Ia were heterogeneous, but also showed that some properties of a SN~Ia did {\\it not} depend on the width of its light curve.  Several studies have shown that SN~Ia spectral properties are related to their photometric properties \\citep[e.g.,][]{Nugent95}.  Some studies examined the correlation between Hubble residuals after light-curve shape correction and spectroscopic properties to produce a smaller Hubble scatter \\citep{Foley08:uv, Bailey09, Blondin11, Chotard11, Nordin11:Si4000}.  However, using this simple approach with only optical spectra provides only minor gains over photometry alone \\citep{Blondin11}.  In a previous study, \\citet[hereafter FK11]{Foley11:vel} showed that the maximum-light intrinsic color of SNe~Ia is highly correlated with their ejecta velocity as probed by the \\ion{Si}{2} $\\lambda 6355$ feature, $v_{\\rm Si~II}$.  Accounting for this correlation both substantially improves the precision and reduces potential biases of SN~Ia distance measurements.  Theoretically, this is understood as increased line blanketing in the higher-velocity SNe, depressing the $B$-band flux relative to the $V$-band flux and causing a redder $B-V$ color.  Another theoretical expectation is that higher-velocity SNe~Ia have more absorption from lower excitation (e.g., \\ion{Fe}{2} vs.\\ \\ion{Fe}{3}) lines, which are distributed such that higher-velocity SNe~Ia should have more absorption in the near-ultraviolet. Nonetheless, not all SN~Ia models show the faster-redder relation found in the data \\citep{Blondin11:2D}.  Although not necessarily the cause of the relation between intrinsic color and $v_{\\rm Si~II}$, the observations are naturally explained by asymmetric explosions \\citepalias{Foley11:vel}, as originally suggested by \\citet{Branch88}.  An asymmetric explosion model was invoked to explain a striking relation between the velocity gradient of the \\ion{Si}{2} $\\lambda 6355$ feature, $\\dot{v}_{\\rm Si~II}$, and the velocity offset of nebular lines at late times \\citep{Maeda10:asym}.  \\citet{Leonard05} found that SNe~Ia with high-velocity features had stronger \\ion{Si}{2} $\\lambda 6355$ line polarization than those without the features.  \\citet{Maund10:asym} expanded upon this work and found a correlation between the amount of polarization in the \\ion{Si}{2} feature and $\\dot{v}_{\\rm Si~II}$, suggesting that $\\dot{v}_{\\rm Si~II}$ is an excellent probe of asymmetry of the outer layers of the SN ejecta.  Since $v_{\\rm Si~II}$ is an excellent proxy for $\\dot{v}_{\\rm Si~II}$ \\citep[hereafter W09]{Wang09:2pop}, one expects $v_{\\rm Si~II}$, $\\dot{v}_{\\rm Si~II}$, intrinsic color at maximum brightness, nebular line velocity offsets, and \\ion{Si}{2} polarization should all be highly correlated. \\citet{Maeda11} confirmed one of these assumptions, showing that nebular line velocity offsets are correlated with maximum-light color. Of all of these observables, $v_{\\rm Si~II}$ is the easiest to obtain at low redshift (and is always obtained if $\\dot{v}_{\\rm Si~II}$ or polarization measurements are made) and is the only observable known to correlate with intrinsic color that we can realistically obtain for high-$z$ SNe.  Besides line velocity, which is determined by the minimum of an absorption feature (corresponding to its maximum absorption), one might expect that the total amount of absorption, as measured by either the full-width at half maximum (FWHM) or pseudo-equivalent width (pEW) of a feature, to strongly correlate with intrinsic color. Both the FWHM and pEW are alternative kinematic probes, measuring the span of the absorbing material in velocity space and may have more physical motivation for a correlation with intrinsic color than absorption velocity.  Although \\ion{Si}{2} $\\lambda 6355$ is a clean, isolated, strong feature in SN~Ia spectra, there is no guarantee that its properties are the most informative for determining intrinsic color. Furthermore, \\ion{Si}{2} $\\lambda 6355$ redshifts out of the optical window at $z \\approx 0.4$.  Therefore, any relation determined with \\ion{Si}{2} $\\lambda 6355$ cannot be applied to most high-$z$ SNe~Ia.  Another strong feature in SN~Ia spectra that has a large velocity distribution and strong velocity evolution with time is \\ion{Ca}{2} H\\&K.  \\citetalias{Foley11:vel} suggested that this feature, which is much bluer than \\ion{Si}{2} $\\lambda 6355$ and can be observed in the optical window up to $z \\approx 1.2$, may be an alternative way to determine intrinsic color. \\citetalias{Foley11:vel} found that the model spectra of \\citet{Kasen07:asym} show a strong relation between intrinsic color and $v_{\\rm Ca~H\\&K}$, similar to what was found with $v_{\\rm Si~II}$.  Current and previous high-$z$ SN~Ia samples typically have a single near-maximum brightness observer-frame optical spectrum per SN \\citep[e.g.,][]{Zheng08, Foley09:year4, Walker11}, making measurements of $\\dot{v}$ impossible.  Many current and future surveys intend to use photometric-only samples of SNe for cosmological studies where only a small subsample of SNe will have spectra.  These samples will also not have the spectral series necessary to determine velocity gradients for high-$z$ SNe~Ia.  Moreover, if intrinsic color can be determined by a single spectrum, telescope time can be used to make this measurement for more SNe than one could if a time series is required.  Since \\citetalias{Foley11:vel} used the \\citetalias{Wang09:2pop} sample, which did not provide spectra or velocity measurements, but only a designation of ``Normal'' or ``High-Velocity'' for SNe, \\citetalias{Foley11:vel} was unable to investigate exactly how intrinsic color depends on ejecta velocity.  Similarly, \\citetalias{Wang09:2pop} did not provide any FWHM or pEW measurements or any measurements for \\ion{Ca}{2} H\\&K.  In this paper, we use the large and homogeneous CfA sample of SN~Ia spectra (Blondin et~al., in prep.), supplemented by literature data (Section~\\ref{s:data}), to construct a homogeneous sample which reduces potential systematic differences, determine a relation between $\\dot{v}$ and $v (t)$ for the two strongest features in a maximum-light SN~Ia spectrum, \\ion{Si}{2} $\\lambda 6355$ and \\ion{Ca}{2} H\\&K (Section~\\ref{s:vgrad}), use the same data to determine a relation between pEW for these features and their gradients (Section~\\ref{s:pew}), use those relations to measure maximum-light values for a large sample of SNe, allowing a direct comparison of these SNe (Sections~\\ref{s:vgrad} and \\ref{s:pew}), and reexamine the relation between ejecta kinematics --- as determined by both the width and velocity of the minimum of the absorption for \\ion{Si}{2} $\\lambda 6355$ and \\ion{Ca}{2} H\\&K --- and intrinsic color (Section~\\ref{s:col}).  We find that the maximum-light velocity of \\ion{Si}{2} $\\lambda 6355$ and \\ion{Ca}{2} H\\&K and the maximum-light pEW of \\ion{Si}{2} $\\lambda 6355$ are correlated with intrinsic color. The intrinsic color scatter is larger for higher-velocity SNe~Ia --- even after removing a linear trend with velocity --- indicating that lower-velocity SNe~Ia are more ``standard crayons.''\\footnote{See \\citet{Peek10} for a description of this nomenclature.}  We discuss these results and conclude in Section~\\ref{s:conc}. ",
        "conclusions": ""
    },
    "1107/1107.1620_arXiv.txt": {
        "abstract": "The sunspot-associated sources at the frequency of 17~GHz give information on plasma parameters in the regions of magnetic field about $B=2000$~G at the level of the chromosphere-corona transition region. The observations of short period (from 1 to 10~minutes) oscillations in sunspots reflect propagation of magnetohydrodynamic (MHD) waves in the magnetic flux tubes of the sunspots. We investigate the oscillation parameters in active regions in connection with their flare activity. We confirm the existence of a link between the oscillation spectrum and flare activity. We find differences in the oscillations between pre-flare and post-flare phases. In particular, we demonstrate a case of powerful three-minute oscillations that start just before the burst. This event is similar to the cases of the precursors investigated by Sych,~R.~\\textit{et al.}~(\\textit{Astron. Astrophys.} \\textbf{505}, 791, 2009). We also found well-defined eight-minute oscillations of microwave emission from sunspot. We interpret our observations in terms of a relationship between MHD waves propagating from sunspot and flare processes. ",
        "introduction": "\\label{S-Introduction}  Quasi-periodic oscillations (QPO) are registered practically in all wavelength ranges and in all structures of the solar atmosphere~(Lites \\textit{et al.}, 1998; Bogdan, 2000; Fludra, 2001; Bogdan and Judge, 2006; Kosovichev, 2009). The periods of oscillations typically range from seconds to hours and perhaps days. The dominant periods of the oscillations are three and five minutes. Most of these oscillations are of intermittent nature: both their amplitude and frequency vary with time. The oscillations are often visible as trains of a few periods. The oscillations with different periods have different physical nature. The short-period oscillations are caused by running waves. Their intermittency is probably caused by instability of physical conditions in the regions of generation and propagation.  Studying oscillations can help us to understand  such fundamental astrophysical problems as accumulation and release of energy, physics of the coronal heating and the origin of flares. The investigation of oscillation processes in the solar corona is a powerful tool for coronal plasma diagnostics~(Nakariakov and Erdelyi, 2009).%  Observations of QPOs of the solar radio emission have been carried out for more than 40 years~\\cite{Durasova71}. The first observations showed that oscillations are mostly produced in solar active regions and their parameters reflect the development of flare activity~\\cite{Kobrin73,Aleshin73}. However, in these studies the spatial and time resolution of radio telescopes was not adequate.  Significant progress in this field was achieved with the help of large modern radio telescopes. Especially rich information was gained from the Nobeyama Radioheliograph (NoRH)~\\cite{Nakajima94}. The NoRH has been observing the Sun since 1992. The observations are carried out every day (7\\,--\\,8 hours daily) with spatial resolution of 10\\,--\\,20~arcsec and cadence of 1 sec.  The most sensitive analysis of the radio oscillations reflecting the  plasma resonance phenomena for MHD oscillations was found for the sunspot associated sources~(Gelfreikh~\\textit{et al.}, 1999; Shibasaki, 2001; Nindos~\\textit{et al.}, 2002; Gelfreikh~\\textit{et al.}, 2006). Such oscillations reflect the oscillations of the plasma parameters (temperature and magnetic field) at the level of the chromosphere-corona transition region (CCTR), where the strength of the magnetic field is about $B=2000$~G. The analysis has shown the presence of regular oscillations of the radio brightness (registered both in $I$ and $V$ Stokes parameters). The periods of oscillations vary from fractions of a minute to several hours. In reality they have a different nature. The nature of oscillations is still in the stage of discussion. However, some mechanisms have been identified. The three- and five-minute oscillations in sunspots are well known from optical observations~\\cite{Uchida75,Horn97} and also analyzed in space-borne data of EUV lines~(Brynildsen~\\textit{et al.}, 1999a; Brynildsen~\\textit{et al.}, 1999b; De~Moortel~\\textit{et al.}, 2002; Brynildsen~\\textit{et al.}, 2003; King~\\textit{et al.}, 2003). It was shown that they usually reflect propagation of MHD waves from the chromosphere towards corona. In some cases the waves propagate along coronal loops rooted in sunspots~% \\cite{De_Moortel02}.  Recently \\inlinecite{Sych09} established the relationship between three-minute oscillations of the microwave emission and quasi-periodic pulsations (QPP) in flares. They found two cases of increase in the power of the three-minute oscillations just before the flares. The oscillations are interpreted as slow magnetoacoustic waves which can cause the QPP.  In many sunspot-associated sources longer periods of QPO (dozens of minutes) are also found. These periods are typical for the coronal loops oscillations~\\cite{Gelfreikh06,Bakunina09,Chorley10}. Some of these may begin in a sunspot region reflecting possibly acoustic oscillations in coronal loops (see also~\\opencite{Nakariakov04}). Much longer periods, up to several hours, are also observed for most of the discussed radio sources. Similar long periods were found from the optical observations of sunspots~(Efremov, Parfinenko, and Solov'ev, 2010). Global oscillations of the Sun may also be reflected in some cases of sunspot radio emission.  In summary, study of QPO observations of  sunspot-associated sources may lead to significant progress in understanding of physical processes in solar plasma. Such observations also contain information on MHD wave propagation in the solar atmosphere. Although so far we have not reached satisfactory interpretation of the physical nature of all observed types of the discussed QPO, there is no doubt that the detailed analysis of their parameters brings new insight concerning the plasma structures in the studied active regions.  The main aim of this work is to study the difference in oscillations observed in pre-flare and post-flare phases and to investigate the reconstruction of the magnetic structure of ARs. ",
        "conclusions": "\\label{S-Discussion}  We have presented three cases that confirm the existence of a link between oscillation spectrum and flare activity.  The analysis of the wavelet spectra of sunspot oscillations at the wavelength of 1.76~cm have confirmed the presence of the variations in the spectra connected with the flare activity. Both disappearance and appearance of some periods were found.  The first case (NOAA 9608, 11 September 2001) shows a decrease of power of the eight-minute oscillations after the first burst and their disappearance after the second, more powerful burst. This case also shows an increase of power of the three- and five-minute oscillations between the bursts. The second case (AR 9866, 14 March 2002) shows the appearance of the eight-minute oscillations after the burst. The most impressive feature in the third case (NOAA 10139, 7 October 2002) is the beginning of well-defined three-minute oscillations 15\\,--\\,20 minutes before the burst.  Two of the three cases demonstrate an increase of power of the short-term oscillations before and during the bursts. This result is in a good agreement with the results obtained by~\\inlinecite{Sych09}. They found similar cases of a gradual increase in the power of the three-minute oscillations before flares in sunspot associated sources. They proposed that slow magnetoacoustic waves propagate from a sunspot along coronal loops upwards the flare site and cause the energy release. This interpretation is  likely to be valid for our cases, too.  In our previous study~\\cite{Abramov11} we showed observational evidence that waves travelled from the chromosphere towards the corona inside the umbral magnetic flux tube of the sunspot. We found similar wave-trains in the oscillation processes at two different levels of the solar atmosphere and measured time shift between wave-trains. In this study we confirmed the relationship between travelling waves and flare activity. One of the possible interpretation of the presented observational material is that the MHD waves propagating from the sunspot can trigger energy release at the flare site.  The  increase of power of the three-minute oscillations before the burst can be used for flare forecasting. Of course, there are some difficulties. First, the detection of the precursor requires a complicated data processing. Second, this effect appears only a few minutes before the flare. So, further studies  are needed. Studies of this effect may shed light on the problem of energy release in solar flares. NoRH is a very suitable instrument for such investigations due to its long series of observations (since 1992), high time (1~sec) and spatial (10~arcsec) resolution for the full solar disk, and  its long daily observing time (up to 8~hours per day).  Our conclusions are as follows:  1. We have confirmed the existence of a link between the oscillation spectrum and flare activity.  2. We found one case of increase of the power of the short-term oscillations before the burst.  3. We interpret our observations in terms of magnetohydrodynamic waves propagating from sunspots and flare processes.  4. The NoRH is a very suitable instrument for such studies.  \\begin{acks} This work was partially supported  by the Presidium of Russian Acade\\-my of Sciences under grant OFN-15.  We are grateful to the anonymous referee and Editors for important remarks, which helped us to improve the paper.  Data used here from Mees Solar Observatory, University of Hawaii, are produced with the support of NASA grant NNG06GE13G.  Wavelet software was provided by C.~Torrence and G.~Compo, and is available at URL: http://paos.colorado.edu/research/wavelets/ \\end{acks}"
    },
    "1107/1107.6000_arXiv.txt": {
        "abstract": "The energy spectrum of cosmic-ray antiprotons ($\\bar{p}$'s) from 0.17 to 3.5 GeV has been measured using 7886 $\\bar{p}$'s detected by BESS-Polar II during a long-duration flight over Antarctica near solar minimum in December 2007 and January 2008. This shows good consistency with secondary $\\bar{p}$ calculations. Cosmologically primary $\\bar{p}$'s have been investigated by comparing measured and calculated $\\bar{p}$ spectra. BESS-Polar II data show no evidence of primary $\\bar{p}$'s from evaporation of primordial black holes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.2415_arXiv.txt": {
        "abstract": "This paper provides testable predictions about starlight polarizations to constrain the geometry of the Galactic magnetic field, in particular the nature of the poloidal component. Galactic dynamo simulations and Galactic dust distributions from the literature are combined with a Stokes radiative transfer model to predict the observed polarizations and position angles of near-infrared starlight, assuming the light is polarized by aligned anisotropic dust grains. S0 and A0 magnetic field models and the role of magnetic pitch angle are all examined. All-sky predictions are made, and particular directions are identified as providing diagnostic power for discriminating among the models. Cumulative distribution functions of the normalized degree of polarization and plots of polarization position angle vs. Galactic latitude are proposed as tools for testing models against observations. ",
        "introduction": "A gap exists between theories of the large-scale Galactic magnetic field and observations that probe the same field. To bridge this gap, predictions from existing Galactic dynamo models that can be tested with current observational techniques, in particular polarization of background starlight, must be generated.  Extensive theoretical work over the past fifty years established the origin and expected nature of galaxy-scale dynamos. \\citet{KZ08} provide a comprehensive review of cosmic magnetic field research and highlight many of the existing questions related to galactic dynamos. A leading theory for magnetic field generation in the Galaxy is the mean-field ``$\\alpha$-$\\Omega$'' disk dynamo \\citep{M78,Ruz88}. This fundamental idea has led to many numerical simulations of Galactic magnetic fields \\citep{E92,B92,B93,FS2000,K02,K03,HWK09,M10}. As these ideas have developed, they became more sophisticated and added much to the understanding of the Galactic magnetic field. Recent advances include understanding the role of magnetic helicity \\citep{BS05,Sh06}, the interplay between the disk and halo fields \\citep{MS08,M10}, and the first global MHD simulation of a cosmic ray-driven dynamo \\citep{HWK09}.  The presence of magnetic fields in the Galaxy has been directly observed via several observational methods: Zeeman splitting of spectral lines \\citep{V68,C93}, Faraday rotation of pulsars \\citep{S68,Man74} and extragalactic sources \\citep{CP62}, and synchrotron emission \\citep{Wester62,Wie62,YZ84,YZ89}. Magnetic fields in the Galaxy have also been indirectly inferred from polarization of background starlight \\citep{HI49,HA49,MF70} and polarized thermal dust emission \\citep{H88}. Zeeman splitting and polarized dust emission only probe high-density regions in interstellar clouds and therefore only weakly contribute to our understanding of the large-scale structure and origin of the Galactic magnetic field. Polarized synchrotron emission has been used to study resolved magnetic fields in other galaxies \\citep[and references therein]{B96} and in the Milky Way, in particular the vertical magnetic fields in the Galactic center \\citep{YZ84}.  Much of the work attempting to understand large-scale magnetic fields in our Galaxy has used Faraday rotation of polarized emission from pulsars and extragalactic sources. Attempts have been made to fit empirical magnetic field models for the Galactic disk to the observed rotation measures \\citep{RK89,MFH08,Sun08,NK10,VE11}. These attempts to measure basic properties of the Galactic magnetic field using rotation measures have proved to be challenging, particularly in light of arguments for an antisymmetric (odd) \\citep{AM88,H97}, or symmetric (even) \\citep{F01,Sun08} axisymmetric field in the disk, and for fields that cannot be so simply classified \\citep{Mao2010}. Additional evidence suggests and that the observed rotation measure distribution may be strongly affected by nearby, small-scale structures \\citep{W2010}.  Polarization of background starlight has previously been used to measure the large-scale orientation of the magnetic field \\citep{H76}, to estimate the curvature and pitch angle (defined as arctan[$B_r/B_\\phi$]) of the field \\citep{H96}, and to study magnetic fields in other galaxies \\citep{Hough96}. However, it has mostly been applied to the study of magnetic fields in and around star forming regions and individual clouds \\citep[e.g.,][]{S85,G95,A98,T06,T07,K07}. By comparison, near-infrared (NIR) polarimetry of background starlight is a relatively untapped resource for the study of the large-scale Galactic magnetic field. Optical polarimetry has typically been restricted to probing magnetic fields within $\\sim$1 kpc \\citep{F02}, limiting its ability to probe the large-scale structure of the magnetic field. However, NIR polarimetry is less extincted by dust and can probe the magnetic field along longer lines of sight. The regular component of the Galactic magnetic field may be better measured by these longer sight lines \\citep[longer than the scale length of the random component;][]{RK89,OS93} as the random magnetic components will tend to null averages. By studying the magnetic field in the outer Galaxy, where the dust density is lower than in the inner Galaxy \\citep{SMB96,DS01}, and where there are fewer supernovae to create magnetic field disturbances \\citep{H87}, NIR polarimetry may be used to probe the quiescent, regular Galactic magnetic field to multi-kpc distances.  Starlight polarimetry is insensitive to the directional polarity of magnetic fields, being sensitive only to their orientations. Therefore, starlight polarimetry cannot directly address questions about the nature of magnetic reversals seen in rotation measure studies. However, starlight polarimetry may shed light on the nature of the poloidal component of the Galactic magnetic field, complementing previous rotation measure studies \\citep{TSS09,Mao2010}. By observing the field projected onto the plane of the sky, this method is sensitive to the Galaxy's poloidal magnetic field component and allows the nature of that field to be measured.  To advance the understanding of the Galactic magnetic field, strong constraints on theory are needed. In this paper, physics-based models of the Galactic dynamo, combined with models of the Galactic dust distribution, are used to make predictions for starlight polarizations that can be tested with current and emerging observational techniques. This will constrain the models, and therefore the physics, producing the Galactic magnetic field. The NIR stellar polarization mechanism and a Stokes radiative transfer model are outlined in \\S 2. Section 3 presents the results of the NIR polarization simulations using the models described in \\S 2. In \\S 4, future comparison of these predictions with NIR polarimetry measurements and the constraints that can be placed on Galactic dynamo models are discussed. This work is summarized in \\S 5. ",
        "conclusions": "The simulations presented here attempt to bridge the gap between theory and observation by making testable predictions for starlight polarization observations based on models of the large-scale structure of the Galactic magnetic field. These predictions, when combined with NIR polarization observations, can be used to test the models for the large-scale magnetic field and their physical underpinnings. Polarization of background starlight is insensitive to magnetic reversals and cannot address the presence or location of reversals. However, these polarizations can allow testing for the existence and nature of the poloidal component of the Galactic magnetic field, since all S0 models yield similar polarization predictions across the sky and all A0 models yield similar predictions that are distinctly different from the S0 models. The predicted longitude shift of the polarization null points with magnetic pitch angle, demonstrated in Models 19-24 (a small shift is seen all of the models, however only Models 19-24 systematically vary the pitch angle), may facilitate an observational test for the Galactic magnetic pitch angle in the disk, similar to the work of \\citet{H96}. This test is fundamentally different from previous attempts to use Faraday rotation of pulsars and extragalactic sources to model the magnetic pitch angle \\citep{RK89,MFH08,NK10,VE11}.  The Galaxy was assumed to have an axisymmetric magnetic field based on theory \\citep{Ruz88,B90,MB92} and observational evidence \\citep{RK89,RL94,Sun08,R10}, but there is some evidence in other galaxies \\citep[particularly in M81,][]{K89,S92} that bisymmetric fields may exist. \\citet{B96} indicate that many of the galaxies showing evidence for bisymmetric magnetic fields also show evidence for galaxy interactions, so we might not expect this magnetic field structure for the Milky Way (based on a lack of major merger events; \\citealt{GWN02}) and the axisymmetric assumption may hold. Higher order azimuthal symmetries might also be possible, but their amplitudes should be relatively small \\citep{B96}.  The measurements needed to distinguish among the various model simulations are well suited to NIR stellar polarimetry. The polarization mechanism used in the simulations works from the NIR through near-UV wavelengths \\citep{SMF75,CM76,W92}. However, NIR light is less attenuated by interstellar dust and can probe magnetic fields along multi-kpc scales, while the shorter wavelengths only probe within about 1 kpc \\citep{F02}. As described by \\citet{SMF75}, the polarization signal in the NIR is weaker than the visible by a factor of four or more. However, NIR polarimetric observations at this level are possible with recent instrumentation \\citep{K06,C07} and should soon provide data able to test models of the large-scale structure of the Galactic magnetic field.  The simplest comparison may be through collecting the polarization behavior exhibited through a set of samples of all Galactic latitudes at a single Galactic longitude, as shown in Figs. \\ref{S0_cut} through \\ref{AA_cut}. To best understand the poloidal component of the Galactic magnetic field, these observations should be made in the outer Galaxy longitude ranges of $\\ell=110-160\\degr$ or $\\ell=200-250\\degr$. These ranges avoid the degenerate regions that reside near $\\ell=0,\\,90,\\,180$, and $270\\degr$ with a $20\\degr$ buffer zone in either direction to allow for possible wander caused by any reasonable magnetic pitch angles (if present). To measure this magnetic pitch angle, a sample of Galactic longitudes should be observed at a single Galactic latitude, say $b=15\\degr$, similar to the approach of \\citet{H96}."
    },
    "1107/1107.5235_arXiv.txt": {
        "abstract": "{We present the first results of the application of the recently found universality of behavior of muon signal \\smu\\ to electromagnetic (EM) signal \\sem\\ ratio with respect to the vertical depth of showers maximum \\xmaxv\\ for mass composition and hadronic interaction studies. Making use of the fact that for zenith angles $>45$ degrees the dependence of \\sigrat\\ on \\xmaxv\\ is very similar for QGSJET~II and EPOS~1.99 we show that this provides the possibility to estimate muon shower content in almost interaction model independent way. To evaluate the excess of signal in the data in respect to Monte-Carlo predictions we propose to use mass independence of the electromagnetic signal. Using the simulations with EPOS~1.99 as a fake data we show that one can determine the absolute scaling factor between these fake data and the interaction model under test (QGSJET~II in our case). Applying this scaling factor to the total and muon signals of QGSJET~II one can make accurate conclusions on the primary mass of samples prepared with EPOS~1.99.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.0297_arXiv.txt": {
        "abstract": "The Kepler Mission is a Discovery mission supported by NASA's Science Mission Directorate, and its primary aim is to discover Earth-like planets in the habitable zone of solar-type stars.  The space telescope was designed with a photometer that monitors the Kepler field in a near continuous manner in order to achieve this goal.  With this mission, the asteroseismology community also benefits from the Kepler data via the abundant time-series photometry.  With a short cadence of 1 minute and long cadence of 30 minute observations, the time coverage for many variable stars is unprecedentedly complete.  The Kepler field also contains the archetype RR Lyr, and the Kepler Asteroseismic Science Consortium (KASC) Working Group for RR Lyrae stars have been working to uncover the mysteries surrounding these stars.  I will provide an overview of the Kepler program in relation to RR Lyrae research. ",
        "introduction": "The Kepler project, NASA's 10th Discovery Mission, was launched in March 2009 with the primary objective of finding Earth-sized planets in the habitable zone of solar-type stars.  While the exoplanet search is ongoing, the impact to asteroseismology projects is enormous.  The high precision time-series photometry provided by Kepler is unprecedented for many areas of variable star research.  For the celebration conference for George W. Preston, it is fitting to present the light curve of the eponym RR Lyr from Kepler.  This proceeding is divided as follows: Section 2 covers the data acquisition of the Kepler data and processing; Section 3 focuses on some published results of work done of RR Lyrae variable stars found in the Kepler field of view as well resources available from the Kepler Guest Observer Office; and Section 4 describes some NASA programs and opportunities to analyze Kepler data. ",
        "conclusions": ""
    },
    "1107/1107.5198_arXiv.txt": {
        "abstract": "{The evolution of dust particles in protoplanetary disks determines many observable and structural properties of the disk such as the spectral energy distribution (SED), the appearance of disks, temperature profile, and chemistry. Dust coagulation is also the first step towards planet formation.} {We investigate dust growth due to settling in a 1D vertical column of a disk. We gradually build up the complexity of the models by considering the effects of porosity, different collision models, turbulence, and different gas models respectively. This way we can distinguish the effects of these physical processes on particle growth and motion. It is known from the 10 micron feature in disk SEDs, that small micron-sized grains are present at the disk atmosphere throughout the lifetime of the disk. We hope to explain such questions as what process can keep the disk atmospheres dusty for the lifetime of the disk and how does the particle properties change as a function of height above the midplane.} {We use a Monte Carlo code to follow the mass and porosity evolution of the particles in time. The used collision model is based on laboratory experiments performed on dust aggregates. As the experiments cannot cover all possible collision scenarios, the largest uncertainty of our model is the necessary extrapolations we had to perform. We simultaneously solve for the particle growth and motion. Particles can move vertically due to settling and turbulent mixing. We assume that the vertical profile of the gas density is fixed in time and only the solid component evolves.} {We find that the used collision model strongly influences the masses and sizes of the particles. The laboratory experiment based collision model greatly reduces the particle sizes compared to models that assume sticking at all collision velocities. We find that a turbulence parameter of $\\alpha = 10^{-2}$ is needed to keep the dust atmospheres dusty, but such strong turbulence can produce only small particles at the midplane which is not favorable for planetesimal formation models. We also see that the particles are larger at the midplane and smaller at the upper layers of the disk. At 3-4 pressure scale heights micron-sized particles are produced. These particle sizes are needed to explain the 10 micron feature of disk SEDs. Turbulence may therefore help to keep small dust particles in the disk atmosphere.} {} ",
        "introduction": "The coagulation of dust particles is the first step towards planet formation. As dust particles are the main source of opacity \\citep{Semenov2003}, they also determine many observable quantities of disks, such as the spectral energy distribution and scattered light images (see below). Due to the vertical component of the stellar gravity, dust particles sediment towards the midplane of the disk. The relative settling velocity between particles of different aerodynamical properties drives coagulation and this process was investigated by several authors \\citep{Safronov1969, Nakagawa1986, Schrapler2004, Dullemond:2004p325, Dullemond2005}.  Observational evidence of the vertical sedimentation of grains exists for a large number of disks, although such evidence is usually indirect \\citep{Henning2011}. Sub-micron grains are present at the disk surfaces as shown by scattered light images in the optical and near infrared (NIR) wavelengths. Such multi-wavelength scattered light images provide evidence for grain growth (\\cite{Watson2007} and references therein). \\cite{Pinte2007} showed by reproducing multi-band images of the binary system of GGTau that the dust scale height for 10 micron-sized particles is roughly half of that for micron-sized particles. The spectral energy distribution (SED) is also affected by settling. \\cite{D'Alessio2006} showed that in order to explain the median SEDs of classical TTauri stars, the dust to gas ratio has to be reduced by a factor of 10 at the disk atmosphere compared to the standard value. There are also indications that the settling of grains is correlated with the age of the disk \\citep{Sicilia-Aguilar2007}. However, the connection between the exact shape of the 10 micron feature of SEDs and sedimentation is not well understood \\citep{Dullemond2008, Juh'asz2010}. \\cite{Apai2005} showed evidence for settling in disks around brown dwarf stars. They concluded that growth, crystallization and settling of dust happens around low mass stars in a similar manner as around intermediate and solar mass stars. This suggests that the first stage of planet formation is a robust process occurring in all kinds of disks.  Sedimentation also affects the vertical temperature structure of the disk. The simulations of \\cite{Aikawa2006} showed that as the dust particles sediment towards the midplane, the opacity is reduced, and the temperature of the gas decreases. As the stellar radiation can now penetrate deeper in the disk, the temperature at intermediate heights increases. The change in the density and temperature structure as well as the radiation field naturally influences the chemistry of the disk atmospheres \\citep{Bergin2007}.  Recently \\cite{Vasyunin2011} investigated the effects of dust growth on disk chemistry and found that the main effect of coagulation on the disk chemical composition comes mostly from sedimentation in their models.  Dust sedimentation not only affects the upper layers of the disk, but also the midplane regions. As long as the dust to gas ratio is much smaller than unity, the dust can be treated as a passive tracer in the gas flow. However, if the dust to gas ratio becomes comparable to unity, the dust will influence the gas. The dust can be accumulated by sedimentation around the midplane of the disk. In such a scenario, a shear is present between the dusty midplane layer and the less dusty upper layers. This shear triggers Kelvin-Helmholtz instability which develops into turbulence. This process was first recognized by \\cite{Weidenschilling1980} and is still under active investigation (e.g. by \\cite{Johansen2006}, \\cite{Chiang2008}, \\cite{Lee2010}).  A collaboration started between the lab-community and the modelers to better constrain the dust evolution in protoplanetary disks, using a realistic collision model that is based on the laboratory experiments. In \\cite{Guttler2010} (henceforth Paper I), we introduced this collision model. In \\cite{Zsom2010} (henceforth Paper II), we used this collision model for the first time in the Monte Carlo (MC) method of \\cite{Zsom2008} (henceforth ZsD08). The models of Paper II were local box models, meaning that the dust evolution was only followed at one location of the disk. These models showed that bouncing plays an important role in dust evolution.  We further develop these models to simulate a 1D vertical column in the disk, thus investigating sedimentation-driven coagulation. We want to better understand the process of sedimentation and the role of particular physical phenomena like porosity of the aggregates, collision models and turbulence.  Previous work by \\cite{Dullemond2005} (henceforth DD05) showed that without a mechanism that reduces the sticking probability of particles in the upper layers of the disk, or without a continuous source of small particles, the observed spectral energy distributions (SED) of TTauri stars should exhibit a very weak IR excess. In contrast, the observed SEDs of TTauri stars have strong IR excesses (e.g. \\cite{Furlan2005, Kessler-Silacci2006, Bouwman2008}). Therefore some grain-retention mechanism is needed to explain the SEDs.  Previous models of grain evolution assumed a continuous cycle of growth and fragmentation, which provides the necessary amount of small particles (e.g. \\cite{Brauer2008a}, \\cite{Birnstiel2009}). However, Paper I and II showed that particle evolution is halted by bouncing and no cycle of growth and fragmentation is present. In this paper we simulate dust evolution driven by Brownian motion, turbulence, and sedimentation in a 1D vertical column of the inner disk and the additional physics of Paper I and II are included. We investigate the time evolution of sedimentation-driven coagulation, and search for ways that can keep a sufficient amount of the small dust particles levitated at several pressure scale heights to explain the observed SEDs of young stars.  The paper is organized as follows. In Sect. \\ref{sec:num} we describe the numerical method used to follow the particle motion and coagulation. We validate the code and increase the complexity of the model step-by-step in Sect. \\ref{sec:intro}. We show the results in Sect. \\ref{sec:res}, finally we discuss those results in Sect. \\ref{sec:disc} and provide a summary in Sect. \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:disc}  \\subsection{Validity of the porosity model} \\label{sec:collmod} The Braunschweig collision model is based on laboratory experiments that used fractal dimension 3 aggregates with an enlargement parameter typically between 7 and 2. As seen in Table \\ref{table:sedi}, the maximum enlargement parameter of the aggregates can be several orders of magnitude larger than the dust aggregates used in the laboratory. More specifically, the hit\\&stick collision regime is followed by a collision type called sticking through surface effects (see Paper I) and bouncing at higher collisional energies. Thus we assume that the transition from fractal aggregates in the hit\\&stick regime to fractal dimension 3 aggregates in bouncing is an instantaneous one. Furthermore, our porosity model for bouncing is calibrated for dust cakes (enlargement parameter of 6), however the enlargement parameter at the end of the hit\\&stick phase is 2-3 orders of magnitude larger than the enlargement parameter of the dust cakes (see Table \\ref{table:sedi}). Therefore, it is questionable whether the porosity model for the restructuring phase is correct. Significant changes in the porosity of aggregates has the potential to significantly alter our collision model and thus the obtained results.    It is still debated how the porosity and the fractal dimension evolves during the evolution of dust aggregates. As this is a central issue in determining the dust properties, we need detailed information on the porosity evolution of dust aggregates. To resolve these issues, we plan to follow the hit\\&stick and restructuring phases of a particle distribution in a self-consistent way by combining the Monte Carlo model of ZsD08 with a molecular dynamics code in a follow-up paper.  \\subsection{Aggregates beyond the snow line} \\label{sec:ice} The particle sizes at the midplane of the disk are rather small ($\\sim$ 100 microns if $\\alpha=10^{-2}$) as we see in the previous section. The question arises: under what conditions could planetesimals form via successive collisions of dust aggregates? Or how do large enough dust aggregates form that could, under favorable conditions, be concentrated in e.g. vortices or turbulent eddies, become self-gravitating, and form eventually planetesimals \\citep{Johansen:2007p65, Cuzzi2008, Youdin2011, Johansen2011}?  One answer to these questions could be icy aggregates. The molecular dynamics simulations of \\cite{Wada2008, Suyama2008} showed that icy aggregates are very resilient to restructuring. They observed sticking between icy aggregates at a relative velocity as high as 50 m/s. The uncertainty in these simulations is the microphysical parameters of the icy monomers (such as critical displacement, surface energy, Young modulus etc.). Recently laboratory experiments were performed using micron-sized ice monomers by \\cite{Gundlach2011}. They measured the rolling energy between icy monomers and it turned out the previously assumed values by \\cite{Wada2008, Suyama2008} are in good agreement with the measured laboratory values. If experiments also confirm that icy aggregates can stick at relative velocities as high as 50 m/s, that would provide a way to form large enough dust aggregates beyond the snow-line in the solar nebula.   \\subsection{Is $\\alpha$ constant as a function of height?} \\label{sec:alpha} \\cite{Gammie1996} proposed the concept of layered accretion disks. If the ionization fraction of the gas is not sufficient to support magneto-rotational instability (MRI - \\cite{Balbus1991}), the turbulence parameter drops and a dead-zone forms at the midplane of the disk. The extent of the dead-zone is uncertain, as the ionization processes of the gas are not well-constrained. For typical TTauri disks it can extend between 0.1 - 4 AU \\citep{D'Alessio1998}. Inside 0.1 AU, the thermal radiation from the star can keep the gas sufficiently ionized for MRI, and outside 4 AU, the gas surface density is typically below 100 g/cm$^2$, therefore cosmic rays can penetrate the disk and keep it MRI active at all heights.  It was proposed by \\cite{Okuzumi2009} that negative charges on the surfaces of grains could prohibit growth. As we use a Monte Carlo code to follow the evolution of aggregates, it is possible to include a third particle property: the amount of charges present on the grain. In order to follow the charge evolution of the grains, we need to solve for the ionization state of the gas (i.e. amount of charges available in the gas phase), how efficiently these charges are collected by the dust grains and how the charges affect the relative velocity of the aggregates.  We plan to investigate how dust evolves in a layered disk model. Small dust particles can very efficiently sweep up charges in the gas. As shown by \\cite{Turner2010}, the dead-zone can extend to 2 $H_g$ for 1 micron-sized particles, but it shrinks below 0.5 $H_g$ for aggregates that are 100 micron in size. In a simulation like the one presented here, this would mean that as the particles grow, the dead-zone shrinks. When the dead-zone disappears, the whole disk becomes MRI active and the particles settled to the midplane might be fragmented and stirred back up. This could lead to initial oscillations before an equilibrium state is reached."
    },
    "1107/1107.2401_arXiv.txt": {
        "abstract": "We investigate the idea that the interaction of Dark Matter (DM) sub-halos with the gaseous disk of galaxies can be the origin for the observed holes and shells found in their neutral hydrogen (HI) distributions. We use high resolution hydrodynamic simulations to show that pure DM sub-halos impacting a galactic disk are not able to produce holes; on the contrary, they result in high density regions in the disk. However, sub-halos containing a small amount of gas (a few percent of the total DM mass of the sub-halo) are able to displace the gas in the disk and form holes and shells. The size and lifetime of these holes depend on the sub-halo gas mass, density and impact velocity. A DM sub-halo of mass  $10^8$ $M_{\\odot}$ and a gas mass fraction of $\\sim 3$$\\%$, is able to create a kpc scale hole, with a lifetime similar to those observed in nearby galaxies. We also register an increase in the star formation rate at the rim of the hole, again in agreement with observations. Even though the properties off these simulated structures resemble those found in observations we find that the number of predicted holes (based on mass and orbital distributions of DM halos derived from cosmological N-body simulations) falls short compared to the observations. Only a handful of holes are produced per Giga year. This leads us to conclude that DM halo impact is not the major channel through which these holes are formed. ",
        "introduction": "Holes and shells have been found in the neutral hydrogen (HI) distribution of many nearby galaxies like M31 (\\citealt{1981A&A....95L...1B}, \\citealt{1986A&A...169...14B}), M33 \\citep{1990A&A...229..362D}, Holmberg II \\citep{1992AJ....103.1841P}, M101 and NGC 6946 \\citep{1993nhns.book.....K}, IC 10 \\citep{1998AJ....116.2363W}, SMC \\citep{1997MNRAS.289..225S} , LMC \\citep{1998ApJ...503..674K} and IC 2574 \\citep{1999AJ....118..273W}. These structures are typically regions of low HI density distributed across the entire galaxy, and come in various sizes ranging from a few 100 pc to about 1.5 kpc. Their expansion velocity sets an upper limit to their dynamical age of about 10 - 60 Myr (cf.  \\cite{1999AJ....118..273W} for IC 2574).   The formation of HI holes and shells has often been ascribed to the action of stellar winds and supernova explosions occurring in OB associations and young stellar clusters (cf. \\cite{1977ApJ...218..377W}, \\cite{1988ARA&A..26..145T}, \\cite{1996ASPC..106...47V}). Based on a simple model of holes created by O and B stars, \\cite{1997MNRAS.289..570O} successfully predicted the observed number distribution of holes in the SMC, lending support to this hypothesis. More recently, the analysis conducted by \\cite{2001AJ....122.3070O}, \\cite{2005AJ....129..160S}, \\cite{2009ApJ...704.1538W} and \\cite{2011ApJ...735...36C} have shown that the stellar content of HI holes in Holmberg I and II, DDO 88 and 165 can release enough mechanical energy to drive the formation of holes on timescales similar to what is observed. In particular, \\cite{2009ApJ...704.1538W} were able to show that all the holes in Holmberg II observed with HST contain multiple stellar populations of different age.  \\cite{2011arXiv1105.4117W} studied  the HI holes in five dwarf galaxies (DDO 181, Holmberg I, M81 Dwarf A, Sextans A and UGC 8508), and failed to detect a young star cluster at the center of the observed HI holes. They thus suggested that large holes may form due to multiple episodes of star formation.  There is more observational evidence against a single burst of star formation being responsible for the creation of HI holes. For example, no robust spatial correlation was found by \\cite{1999AJ....118.2797K} between the distribution of the H$\\alpha$ emission and the HI holes in the LMC. \\cite{2005MNRAS.360.1171H} estimated that the holes in the SMC not associated with a young star cluster are a factor of 1.5 more than those exhibiting relatively young stars at their center. In the Galaxy, the distribution of radio pulsars compared to that of holes seem to indicate that holes can not be produced by the supernova explosions of a single age stellar population \\citep{2004ApJ...606..326P}. \\cite{2008ApJ...687.1004P} found that the most recent episode of star formation has taken place preferentially at the rims of the HI holes in IC 2574, exception made for the supergiant shell which indeed embeds a young star cluster (cf. \\cite{2000AJ....120.1794S}, \\cite{2009ApJ...691L..59W}). \\cite{1999AJ....118..323R} observed the HI holes in Holmberg II in search of the descendants (stars of spectral types B, A and F) of those clusters whose supernovae would have produced the holes. They measured an integrated light of the stars within the holes inconsistent with the hypothesis of a young star cluster triggering the formation of a HI hole. In the case of IC 1613, its largest HI shell was found to host about 27 OB associations whose energy can not sustain the formation and expansion of the shell (\\cite{2004A&A...413..889B}, \\cite{2006A&A...448..123S}).     Some authors have proposed alternative formation hypotheses to the SNe origin. \\cite{1998ApJ...501L.163E} and \\cite{1998ApJ...503L..35L} postulate that a high-energy gamma ray burst (GRB) from the death of a single massive star could create kpc sized holes in the Inter Stellar Medium (ISM), thus offering an explanation for holes without a detectable underlying cluster. These authors assume the energy from GRBs is emitted isotropically. However, GRBs release most of their energy in bi-directional beams (eg., \\citealt{1977MNRAS.179..433B} mechanism), making this scenario less likely to produce large HI holes.   Another mechanism proposed is the infall of gas clouds \\citep{1987A&A...179..219T}. One observational prediction of this model is a half-circle arc seen in an HI position-velocity diagram. The half-circle arc arises from the gas being pushed to one direction, corresponding to the direction of travel of the high velocity clouds. Some observational support for this idea is reported by Heiles (\\citealt{1979PASP...91Q.611H}, \\citealt{1984ApJS...55..585H}) who point out that the most energetic Galactic shells in their study all have half-circle arc signatures in position-velocity space ( \\citealt{1991A&A...244L..29K}).   In this paper we investigate an alternative formation scenario, where holes and shells in extended HI disks are the result of interaction of the gaseous disk with dark matter (DM) sub-halos. The Lambda Cold Dark Matter (LCDM) model predicts the existence of thousands of dark matter substructures within the dark matter halo of every galaxy (e.g. \\citealt{2008MNRAS.391.1685S} and reference therein). The majority of these sub-halos is ``dark'' i.e. does not host a satellite galaxy or any visible stellar structures (e.g. \\citealt{2010MNRAS.402.1995M}, \\citealt{2011arXiv1102.2526F}). If these sub-halo population is able to produce the observed holes in galactic HI disks, like the one of IC 2574, then this might provide a new way of detecting the presence of non-luminous DM halos. For example, the number of holes could then be linked statistically to the amount of substructure in the dark matter distribution  and thus provide a unique way to asses the nature of dark matter and by extension to test the $\\Lambda$CDM model.   Perturbations on a stellar disk due to (massive) satellites have been extensively studied in the recent years (e.g. \\citealt{2008ApJ...688..254K}, \\citealt{2010MNRAS.403.1009M} and references therein), while less attention has been given to the effect of low mass DM clumps on an extended gaseous disk, with few exceptions. \\cite{2006ApJ...637L..97B}  investigate how the impact of dark matter sub-halos orbiting a gas-rich disk galaxy embedded in a massive dark matter halo influences the dynamical evolution of the outer HI gas disk of the galaxy. They show that the impact of dark matter sub-halos (``dark impact'') can be important for better understanding the origin of star formation discovered in the extreme outer regions of disk galaxies.  They also discuss the possibility that this kind of dark impact will be able to produce holes in the gas distribution. In their study they adopted a model, which did not include multiple events, a live dark matter halo (they adopted an analytic fixed potential), cooling and star formation.  A new attempt to detect the imprint of the dark satellites in the HI disk of the Milky Way has been recently made. In a series of papers \\cite{2009MNRAS.399L.118C} and \\cite{2011ApJ...731...40C} examine tidal interactions between perturbing dark sub-halos and the gas disk of the Milky Way using high-resolution Smoothed Particle Hydrodynamics simulations. They compare their numerical results to the observed HI map of the Milky Way, and find that the Fourier amplitudes of the planar disturbances are best fitted by a perturbing dark sub-halo with a mass that is one-hundredth of the Milky Way with a pericentric distance of 5 kpc. More recently  ( \\citealt{2011arXiv1102.3436C}) develop a perturbative approach to study the excitation of disturbance in the extended atomic hydrogen discs of galaxies produced by  passing dark matter sub-halos. They show that the properties of dark matter sub-halos can be inferred from the profile and amplitude of the different perturbed energy modes of the disk.   Motivated by these recent studies, in this paper we use high resolution hydro-dynamical simulations to study in details the interaction of DM sub-halos with a galactic gaseous disk. Our primary goal is to see under which conditions DM satellites are able to create holes that resemble the observed ones and whether the majority of these holes can be explained by dark satellites-disk interactions. We model our primary galaxy on the nearby dwarf galaxy IC 2574 and try to replicate its observed features in our simulations. After having described our numerical setup, we will first present results of a single satellite-disk interaction, for different satellite orbital parameters and gas content. Then we will use satellite parameters (mass, velocity and position) directly obtained from high resolution N-body simulations to study the frequency of DM sub-halo disk encounters. Finally we will discuss the implications of our results. ",
        "conclusions": "Atomic hydrogen (HI) observations of nearby galaxies reveal complex gas distributions. In many systems, the neutral interstellar medium (ISM) contains holes, shells, and/or cavities. In order to understand the origin of these features, we numerically investigate how the impact of dark matter sub-halos orbiting a gas-rich disk galaxy embedded in a massive dark matter halo influences the dynamical evolution of the HI gas disk of the galaxy. We mainly focus our attention on the creation of large holes ($R>1$ kpc) in the HI distribution, commonly found in observations of nearby galaxies.  We create a central galaxy resembling the properties of the well-studied dwarf galaxy IC2475 and we bombard its gaseous disk with dark matter dominated satellites. The dark matter, stellar and gas components of both the primary galaxy and the satellites are all live (i.e. made of particles) and a full hydro-dynamical code ({\\sc gadget2}) including cooling, star formation and feedback is used for these merger simulations.  Our experiment shows that a pure dark matter sub-halo ($M=10^8 M_{\\odot}$) is not able to displace the gas in the disk, instead due to gravitational focusing it gives rise to a localized high density region. This kind of dark impact might be able to induce star formation in the outer part of the extended gas disk, but cannot create large holes. This result was not unexpected: DM particles interact only through gravity and do not have the necessary contact force to push the gas away from the disk.  These high density tidal imprints can be characterized and studied in order to obtain the properties of the impacting dark sub-halo as done by \\cite{2011arXiv1102.3436C} especially for relatively massive dark satellites (1:100).  To produce holes we need particles which have contact forces, hence we assume a small amount of gas to be present in the dark matter sub-halo. Even a gas fraction as low as $0.8\\%$ in a halo of mass $10^8 M_{\\odot}$ with a velocity of 150 km/s is able to produce a detectable low density region. On average holes have a lifetime of about $20-60$ Myr, depending on the halo density, gas mass and impact velocity.  We then use satellite properties directly extracted from high resolution Nbody cosmological simulation (Via Lactea II, Diemand \\etal 2008), to check how many holes are predicted in the commonly assumed Cold Dark Matter model.  These cosmological motivated runs show that sub-halos with a relatively high gas fraction (5\\%), are able to produce a total of about 3-4 large holes ($R>1$ kpc) in an integration time of $1$ Gyr. This number is significantly lower than the number of observed holes of the same size in IC2475 galaxy. If the effect of ram pressure is taken into account, the dark matter satellites tend to loose a significant fraction of their gas content, making them even less efficient at perturbing the HI disc.  We conclude that, although DM matter satellites with a modest gas content are in principle able to create holes with a radius of order 1 kpc, disk - sub-halo interaction are not the primary channel  through which these holes form in real galaxies. We consider it likely that other astrophysical processes like supernova explosions, stellar winds, high velocity clouds, or a combination of the above factors as the main causes for the observed complex geometry of extended HI disks."
    },
    "1107/1107.5826_arXiv.txt": {
        "abstract": "Icy grain mantles are commonly observed through infrared spectroscopy toward dense clouds, cloud cores, protostellar envelopes and protoplanetary disks. Up to 80\\% of the available oxygen, carbon and nitrogen are found in such ices; the most common ice constituents -- H$_2$O, CO$_2$ and CO -- are second in abundance only to H$_2$ in many star forming regions. In addition to being a molecular reservoir, ice chemistry is responsible for much of the chemical evolution from H$_2$O to complex, prebiotic molecules. Combining the exisiting ISO, Spitzer, VLT and Keck ice data results in a large sample of ice sources ($\\sim$80) that span all stages of star formation and a large range of protostellar luminosities ($<$0.1--10$^5$ L$_\\odot$). Here we summarize the different techniques that have been applied to mine this ice data set on information on typical ice compositions in different environments and what this implies about how ices form and evolve during star and planet formation. The focus is on how to maximize the use of empirical constraints from ice observations, followed by the application of information from experiments and models. This strategy is used to identify ice bands and to constrain which ices form early during cloud formation, which form later in the prestellar core and which require protostellar heat and/or UV radiation to form. The utility of statistical tests, survival analysis and ice maps is highlighted; the latter directly reveals that the prestellar ice formation takes place in two phases, associated with H$_2$O and CO ice formation, respectively, and that most protostellar ice variation can be explained by differences in the prestellar CO ice formation stage. Finally, special attention is paid to the difficulty of observing complex ices directly and how gas observations, experiments and models help in constraining this ice chemistry stage. ",
        "introduction": "\\noindent The evolution from starless to star-forming cores and further on to (extra)solar systems is accompanied by a rich chemistry, which will affect planet and planetesimal compositions. Much of this chemical evolution takes place on interstellar grain surfaces, in icy mantles. Ices are generally observed in the cold and dense interstellar medium that are the nurseries of stars and are also abundant in our own solar system in comets and other outer solar system bodies. Such icy bodies are proposed to have delivered volatiles to the young earth. Understanding their link to the chemistry observed during star formation is key to constrain the amount of organic material delivered to Earth as well as to exo-planets.  The first ices were detected in the interstellar medium almost 40 years ago by \\cite{Gillet73}.  H$_2$O and CO were established early on as common ice constituents with abundances reaching 10$^{-4}$ $n_{\\rm H_2}$. This makes ices the most common molecules after H$_2$ in many star forming regions. CO$_2$ is the third major ice constituent. Because of abundant atmospheric CO$_2$, it was only detected in the ISM after the launch of IRAS by \\cite{dHendecourt89}. Following IRAS, the {\\it Infrared Space Observatory} (ISO) and the {\\it Spitzer Space Telescope} have, together with ground-based efforts, continued the ice exploration and interstellar ices. In addition to the major ice constituents, H$_2$O, CO and CO$_2$  are observed to often contain smaller amounts of CH$_3$OH, CH$_4$, NH$_3$ and OCN$^{-}$/XCN  (see \\cite{Boogert04} for a review).  The simple ices observed in the ISM may become the source of complex molecules as described by \\cite{Charnley92} and \\cite{Garrod08} and determining ice abundances and production channels in star forming regions is of great  interest for studies of prebiotic chemistry. Ices can form from direct freeze-out of molecules from the gas phase, through atomic addition on grain and ice surfaces, and through energetic processing of already existing ices. At the low temperatures ($<$20~K) prevalent in cloud cores, atoms and molecules will stick to the grain upon impact. In this stage grains are well protected from external UV radiation and atomic addition reactions probably dominate. \\cite{Tielens82} worked out in detail how the atoms and molecules can be hydrogenated and oxygenated to form all commonly observed ices, except for CO, which condenses directly from the gas-phase. In this scheme, H$_2$O forms from O+H, CH$_3$OH from CO+H etc. CH$_3$OH can also form from UV processing of H$_2$O:CH$_4$ ice mixtures, however, and there has been some debate on which formation pathway dominates. \\cite{Gibb04} and others have also suggested that at least some CO$_2$ and XCN form from thermal or UV processing, mainly based on high-mass protostellar studies with {\\it ISO}.  {\\it Spitzer} and complimentary ground-based surveys have more recently allowed us to observe ices toward large numbers of low-mass protostars and cloud cores, which have provided new observational constraints on how ice formation and evolution depend on the local environment. The aim of this contribution is to review the constraints supplied by ice observations toward both {\\it ISO} and {\\it Spitzer} targets, on the identifications of ice bands and of the main formation paths of the identified carriers. The figures are taken from \\\"Oberg et al. (2011, subm. to ApJ) unless otherwise noted. ",
        "conclusions": "Thanks to a combination of ground-based and space telescope ice observations and years of laboratory ice spectroscopy and ice chemistry experiments we have good constraints on the typical ice abundances and ice compositions in a range of interstellar and circumstellar environments. A key result coming out of these observations is that most ices up to CH$_3$OH in complexity have already formed by the time the protostar turns on. Ice abundance variations must therefore be mainly due to variations in the prestellar ice formation processes. Variations in CO freeze-out and the subsequent CO-based chemistry is a likely main contributor to the observed differences in ice abundances in different lines of sight, but other processes may be important as well. Thermal and UV processing of ices by the protostar is important for the continued evolution of the ice, leading up to the formation of hot core molecules and the complex organics found in comets. To further constrain the ice evolution from cloud cores to comets will require a combination of high-resolution IR studies to detect complex molecules directly in the ice, millimeter observations toward protostars and disks to image desorbed ice chemistry products, and more detailed experiments and models of both the early ice formation stages and of the reactions that convert simple ices into more complex ones."
    },
    "1107/1107.1258_arXiv.txt": {
        "abstract": "Space reddenings are derived for 15 Galactic Cepheids from dereddening CCD {\\it BV(RI)}$_C$ data for AF-type stars in the immediate vicinities of the variables, in conjunction with 2MASS reddenings for BAF-type stars in the same fields. Potential reddening solutions were analyzed using the variable-extinction method to identify stars sharing potentially similar distances and reddenings to the Cepheids, several of which have large color excesses. The intrinsic {\\it BV(RI)}$_C$ color relation for AF dwarfs was modified slightly in the analysis in order to describe better the colors observed for unreddened stars in the samples. ",
        "introduction": "} Because of the deleterious effects of interstellar reddening and extinction, it is difficult to establish an empirical picture of the Cepheid instability strip based entirely upon observations of Milky Way Cepheids. The use of extragalactic Cepheids is not necessarily a practical solution to the problem, since the available data for Milky Way Cepheids are generally more extensive and of greater precision, and the effects of internal reddening within other galaxies are not as well established as they are locally. Such considerations justify the continued use of the Galactic sample in studies aimed at establishing reliable intrinsic properties of classical Cepheid variables.  The determination of accurate reddenings for Cepheids has traditionally followed three different routes: (i) from field reddenings of specific objects (e.g., binary, cluster, association, or isolated Cepheids) based upon the analysis of photometric data for early-type stars sharing the same lines of sight, (ii) from observations mainly of bright Cepheids involving photometric parameters designed to be independent of interstellar reddening, and (iii) from using standard reddening laws and a calibrated intrinsic color relation (either observational or model generated) to deredden photometric observations for large samples of Cepheids. Methods (i) and (ii) are the most reliable means of establishing reddenings for individual Cepheids since method (iii) may entail use of period-color relations \\citep[e.g.,][]{fe90a,fe90b}, which do not account for the intrinsic spread in effective temperature of the Cepheid instability strip and may generate erroneous results \\citep[see][]{tu95}.  For method (ii), spectroscopic indices related to stellar effective temperature and designed to be independent, or relatively independent, of interstellar and atmospheric extinction include the $\\Gamma$-index \\citep{kr60,sj89}, which measures the depression of the {\\it G}-band of CH ($\\lambda$4305) relative to the local continuum, the $\\beta$-index, which samples the strength of the H$\\beta$ Balmer line of hydrogen relative to the surrounding continuum, and the {\\it KHG}-index of Brigham Young University \\citep{mp69,mc70,fm72}, which measures the strengths of Ca {\\sc II} K ($\\lambda$3933), Balmer H$\\delta$ ($\\lambda$4101), and the {\\it G}-band through narrow band interference filters in a manner independent of extinction. A more recent technique \\citep{sl90} involves the use of line depth ratios between C {\\sc I} and Si {\\sc II} lying on the Brackett continuum near 10728 \\AA. A similar technique was used by \\citet{kr98} using spectral lines falling in the red region of Cepheid spectra. \\citet{ko08} have taken the technique to its ultimate level of sophistication by using high resolution optical spectra of Cepheids throughout their cycles in conjunction with stellar atmosphere models to track their changes in effective temperature, and thus intrinsic broad band color.  The use of close neighbours to deduce reddenings for Cepheids appears to date from its application to the blue companion of $\\delta$ Cep by \\citet{eg51}; it has also been used recently for Cepheids studied with the Hubble Space Telescope \\citep{be07}. But the existing sample includes only $\\sim40$ Cepheids of known space reddening \\citep{lc07}. This study presents new space reddenings for 15 Galactic Cepheids derived from CCD {\\it BV(RI)}$_C$ photometry for stars in the immediate fields of the variables. The goal was to test the feasibility of deriving Cepheid reddenings based upon only {\\it BV(RI)}$_C$ observations for stars in the fields of the Cepheids, as well as to augment the limited sample of Cepheids with space reddenings. As demonstrated here, the methodology appears to be provide useful results that extend the sample of Cepheids with independently derived reddenings. ",
        "conclusions": "} The present study was undertaken as a feasibility test of the use of CCD {\\it BV(RI)}$_C$ photometry for studying the space reddening of Galactic Cepheids. The {\\it BV(RI)}$_C$ data by themselves were sometimes insufficient to define the field reddenings of each Cepheid unambiguously, but the addition of 2MASS reddenings and initial ``predictions'' guided the process successfully. The number of Galactic Cepheids of well-established field reddening has been increased from 40 \\citep{lc07} to 55 by this study, and the sample now includes a few Type II Cepheids, some of which appear to display slightly bluer intrinsic {\\it (V--R)}$_C$ and {\\it (V--I)}$_C$ colors relative to their classical Cepheid cousins. A comparable test using reddenings derived from 2MASS {\\it JHK}$_s$ colors indicates that the 2MASS survey may also be suitable for the derivation of space reddenings, for a variety of stellar types and not just Cepheids. Further studies of that possibility have already been initiated \\citep*{me08a,me08b}.  The results of this study have already been used in a preliminary mapping of the Cepheid instability strip using Cepheid reddenings tied to space reddenings and spectroscopic reddenings of Galactic Cepheids \\citep{tu01}, the results of which can be seen in the data plotted in Fig.~\\ref{fig6}. The addition of {\\it U}-band observations would have assisted the analysis considerably, but it is important to note that accurate Johnson system {\\it U}-band observations are difficult to obtain with some CCD/filter combinations \\citep[see comments by][]{tu11}. Despite that, the success of the present study using both {\\it BV(RI)}$_C$ and {\\it JHK}$_s$ photometry should spur further investigations that take advantage of existing survey photometry to study the reddening in interesting Galactic star fields.  \\subsection*{{\\rm \\scriptsize ACKNOWLEDGEMENTS}} \\small{The present study was supported in part by research funding awarded through the Natural Sciences and Engineering Research Council of Canada (NSERC), through the Russian Foundation for Basic Research (RFBR), and through the program of Support for Leading Scientific Schools of Russia. 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.}"
    },
    "1107/1107.3777_arXiv.txt": {
        "abstract": "{As the mid-IR luminosity represents a good isotropic proxy of the AGN power, a low X-ray to mid-IR luminosity ratio is often claimed to be a reliable indicator for selecting Compton-thick AGN. We assess the efficiency of the X-ray to mid-IR luminosity ratio diagnostic by examining the 12\\,\\micron {\\it IRAS} AGN sample (intrinsic $\\rm L_{2-10 keV}>10^{42}$\\,\\lunits) for which high signal-to-noise \\xmm observations have been recently become available. We find that the vast majority (ten out of eleven) of the AGN that have been classified as Compton-thick on the basis the X-ray spectroscopy by Brightman \\& Nandra present a low \\lxl6 luminosity ratio, i.e. lower than a few percent of the average AGN ratio which is typical of reflection-dominated Compton-thick sources. We note that at low \\lxl6 ratios we also find a comparable number of AGN, most of which are heavily absorbed but not Compton-thick. This implies that although most Compton-thick AGN present low \\lxl6 ratios, at least in the local, Universe, the opposite is not necessarily true. Next, we extend our analysis to higher redshifts. We perform the same analysis in the Chandra Deep Field South where excellent quality \\chandra (4\\,Ms) and \\xmm (3\\,Ms) X-ray spectra are available. We derive accurate X-ray luminosities for  \\chandra sources using X-ray spectral fits, as well as 6\\micron luminosities from spectral energy distribution fits. We find eight AGN (intrinsic $\\rm L_{2-10\\,keV}>10^{42}$\\,\\lunits) with low \\lxl6 ratios in total, after excluding one source where the 6\\,\\micron emission primarily comes from star-formation. One of these sources has been already demonstrated to host a Compton-thick nucleus, while for another one at a redshift of $z=1.22$ we argue it is most likely Compton-thick on the basis of its combined \\chandra and \\xmm spectrum. In agreement with the low redshift sample, we find a large number of non Compton-thick `contaminants' with low X-ray to mid-IR luminosity ratios. The above suggest that a low \\lxl6 ratio alone cannot ascertain the presence of a Compton-thick AGN, albeit the majority of low \\lxl6 AGN are heavily obscured. More interestingly, the two most reliable Compton-thick AGN in the high redshift Universe have high \\lxl6 ratios, showing that this method cannot provide complete Compton-thick AGN samples. \\keywords {X-rays: general; X-rays: diffuse background; X-rays: galaxies; Infrared: galaxies}} ",
        "introduction": "The deep X-ray Universe has been probed at unparalleled depth thanks to the \\chandra mission. The deepest \\chandra X-ray surveys revealed an AGN sky density above 10.000$\\rm deg^{-2}$ \\citep{Alexander2003,Luo2008,Xue2011}, probing flux limit of $f_{2-10\\,keV}\\approx 1 \\times 10^{-16}$\\,\\funits \\citep[for a review see][]{Brandt2005}. In contrast, the optical surveys for QSOs reach sky densities of about few hundred per square degree at a magnitude limit of $B=22$mag \\citep[e.g.][]{Wolf2003}. Optical spectroscopic methods, which detect AGN using the [OIII] ($\\rm \\lambda 5007 \\AA$) or the [NeV] lines ($\\rm \\lambda 3427 \\AA$), \\citep{Hao2005,Bongiorno2010,Gilli2010}, manage to recover a large fraction of low-luminosity AGN \\citep{Georgantopoulos2010} but these are applicable only at a limited redshift range, because these lines are redshifted outside the optical band at high redshifts. Therefore, hard (2-10\\,keV) X-ray surveys have provided so far the most powerful method for the detection of AGN and thus have provided the most detailed picture of the accretion history of the Universe \\citep[e.g.][]{Ueda2003, LaFranca2005, Aird2010}.  However, even the efficient hard X-ray surveys may be missing a substantial fraction of the most heavily obscured sources, the Compton-thick AGN, which have column densities $\\rm >10^{24}\\,cm^{-2}$. At these high column densities, the attenuation of X-rays is mainly due to Compton-scattering rather than photoelectric absorption. Although there are only a few dozen of Compton-thick AGN identified in the local Universe, \\citep[see][for a review]{Comastri2004}, there is concrete evidence for the presence of a much higher number. The peak of the X-ray background at 20-30\\,keV \\citep[e.g.][]{Frontera2007,Churazov2007,Moretti2009} can be reproduced by invoking a significant number of Compton-thick sources at moderate redshifts. However, the exact density of Compton-thick sources required by X-ray background synthesis models still remains an open issue \\citep*{Gilli2007,Sazonov2008,Treister2009}. Additional evidence for the presence of a numerous Compton-thick population comes from the directly measured space density of black holes in the local Universe \\citep[see][]{Soltan1982}. It is found that this space density is a factor of 1.5-2 higher than that predicted from the X-ray luminosity function \\citep{Marconi2004,Merloni2008}, although the exact number depends on the assumed efficiency in the conversion of gravitational energy to radiation.  The {\\it INTEGRAL} and {\\it SWIFT} missions explored the X-ray sky at energies above 10\\,keV with the aim to provide the most unbiased samples of Compton-thick AGN in the local Universe. Owing to the limited imaging capabilities of these missions (coded-mask detectors), the flux limit probed is very high ($\\sim 10^{-11}$\\,\\funits), allowing only the detection of AGN at very low redshifts. These surveys did not detect large numbers of Compton-thick sources \\citep[e.g.][]{Ajello2008,Tueller2008,Paltani2008,Winter2009,Burlon2011}. The fraction of Compton-thick AGN in these surveys does not exceed a few percent of the total AGN population. In contrast, optical and mid-IR surveys yield Compton-thick AGN fractions which range between 10 and 20\\% \\citep{Akylas2009,Brightman2011}. As \\citet{Burlon2011} point out, it is possible that even these ultra-hard surveys are biased against the most heavily obscured  ($>2\\times 10^{24}$\\,\\cunits) Compton-thick sources.  At higher redshifts, a number of efforts have been made to identify Compton-thick AGN by means of either X-ray hardness ratios \\citep{Brunner2008,Gilli2011} or X-ray spectroscopy \\citep{Tozzi2006,Georgantopoulos2009,Comastri2011,Feruglio2011} in the deepest \\chandra and \\xmm observations. In particular, \\citet{Comastri2011}, provided the most direct X-ray spectroscopic evidence yet for the presence of Compton-thick nuclei at high redshift, reliably identifying two Compton-thick AGN at z=1.536 and z=3.700.  The mid-IR wavelengths have attracted much attention for providing an alternative way to detect heavily obscured systems. This is because the absorbed radiation by circumnuclear dust is re-emitted in the IR part of the spectrum \\citep*{Soifer2008}, rendering heavily obscured AGN copious mid-IR emitters. In such systems, the 2-10\\,keV X-ray emission can be diminished by almost two orders of magnitude \\citep[e.g.][]{Matt2004}, while at the same time the isotropic mid-IR emission remains largely unattenuated. \\citet{Daddi2007,Fiore2008,Georgantopoulos2008,Treister2009b,Eckart2010,Donley2010} propose that a fraction of infrared excess, $\\rm 24\\,\\mu m$-bright sources which are undetected in X-rays are associated with Compton-thick AGN. Their main argument is based again on the fact that the X-ray to IR luminosity ratio is very low, suggestive of significant amounts of absorption \\citep[but see][]{Georgakakis2010,Georgantopoulos2011}.  The detection of a low \\lxl6 luminosity ratio has recently been used as the main instrument for the detection of faint Compton-thick AGN which remain undetected in X-rays \\citep[e.g.][]{Alexander2008,Goulding2011}. This is because the $\\rm 6\\,\\mu m$ luminosity is an excellent proxy of the AGN power, as it should be dominated by very hot dust which is heated by the AGN \\citep[e.g.][]{Lutz2004,Maiolino2007}. At these wavelengths the contribution by the stellar-light and colder dust heated by young stars should be minimal. \\citet{Yaqoob2011} criticise the above results claiming that the \\lxl6 ratio is not a robust indicator of Compton-thick obscuration. They perform Monte-Carlo X-ray spectral simulations of heavily obscured AGN. They caution that Compton-thick AGN may have the same X-ray to mid-IR luminosity ratio with heavily obscured AGN which have a column density of a few times $10^{23}$\\,\\cunits, depending on the exact value of the covering fraction and photon index.  In this paper, we attempt to assess the efficiency of the \\lxl6 ratio method, i.e. to estimate the percentage of Compton-thick AGN among low \\lxl6 sources, but also to evaluate whether there are any bona-fide Compton-thick AGN with substantially high \\lxl6 ratios. First, we are using local AGN, the IRAS 12\\,\\micron sample of \\citet*{Rush1993}, to assess the reliability and effectiveness of the \\lxl6 ratio selection. For this sample there are excellent quality X-ray spectroscopic observations available \\citep{Brightman2011}, and therefore we know a priori which objects are Compton-thick. Second, we are selecting candidate Compton-thick AGN among the low \\lxl6 AGN in the CDFS. We are presenting a detailed X-ray spectral analysis of these sources, using the most sensitive X-ray data ever obtained, the 4\\,Ms \\chandra data, combined with the 3\\,Ms \\xmm data. Our aim is to detect unambiguous signs of Compton-thick obscuration, such as direct detection of a large column density, or a flat spectral index, or a high equivalent-width Fe\\,K$\\alpha$ line. We adopt $\\rm H_o=75\\,km\\,s^{-1}\\,Mpc^{-1}$, $\\rm\\Omega_{M}=0.3$ and $\\Omega_\\Lambda=0.7$ throughout the paper. ",
        "conclusions": "In this paper, we are assessing the efficiency of the low $\\rm L_X/L_{6\\,\\mu m}$ ratio technique for discovering Compton-thick AGN. The principle behind this method is that in Compton-thick sources, the X-ray rest-frame luminosity, is diminished by almost two orders of magnitude while the 6\\,\\micron luminosity is left unattenuated.  The advantages of the present work in comparison with previou applications of the \\lxl6 method are the following: \\begin{itemize} \\item{We use a 'calibration' sample in the local Universe. This consists of a sample with high-quality \\xmm observations available \\citep{Brightman2011b} and thus with a priori knowledge of the fraction of Compton-thick sources.} \\item{At higher redshift, we employ the most sensitive X-ray observations so far available, namely the 4\\,Ms \\chandra observations, combined with the 3\\,Ms \\xmm observations in the CDFS \\citep{Comastri2011}.} \\item{We derive the X-ray luminosities using proper X-ray spectral fits instead of the usually employed X-ray fluxes in combination with an average X-ray spectral index.} \\item{We use SED decomposition to ascertain that the nuclear contribution is not contaminated by significant amounts of star-forming emission.} \\end{itemize}   Our conclusions can be summarised as follows: \\begin{itemize} \\item{In the local sample the vast majority (ten out of eleven) of the Compton-thick AGN appear to have low \\lxl6 ratios. However, the opposite is not true, i.e. there is a large number of low \\lxl6 sources which are probably not Compton-thick. This would make the efficiency of the \\lxl6 method for selecting Compton-thick sources 50\\%.} \\item{In the CDFS we select eight low \\lxl6 AGN with intrinsic luminosities $\\rm L_X>10^{42}$\\,\\lunits (after exclusion of one source where the 6$\\rm \\mu m$ emission comes from star-formation). One of these (CDFS-309 at a redshift of $z\\approx2.6$) is already shown to host a Compton-thick nucleus, while for another one (CDFS-398 at z=1.222) we argue here for the first time that it is Compton-thick.} \\item{A large fraction of the low \\lxl6 CDFS sources cannot be confirmed as Compton-thick on the basis of the X-ray spectroscopy, but they are nevertheless highly obscured, having column densities above $10^{23}$\\,\\cunits. This suggests that the \\lxl6 ratio cannot be used on its own to reliably classify sources as Compton-thick.} \\item{Finally, there at least two bona-fide Compton-thick sources from \\citet{Comastri2011} which do not appear to present low \\lxl6 ratios, casting further doubt on the validity of this method for selecting Compton-thick sources.} \\end{itemize}"
    },
    "1107/1107.4361_arXiv.txt": {
        "abstract": "We present a study of dense molecular gas kinematics in seventeen nearby protostellar systems using single-dish and interferometric molecular line observations. The non-axisymmetric envelopes around a sample of Class 0/I protostars were mapped in the \\nthp\\ ($J=1\\rightarrow0$) tracer with the IRAM 30m, CARMA and PdBI as well as \\nht\\ (1,1) with the VLA. The molecular line emission is used to construct line-center velocity and linewidth maps for all sources to examine the kinematic structure in the envelopes on spatial scales from 0.1 pc to $\\sim$1000 AU. The direction of the large-scale velocity gradients from single-dish mapping is within 45$^{\\circ}$ of normal to the outflow axis in more than half the sample. Furthermore, the velocity gradients are often quite substantial, the average being $\\sim$2.3 \\kmspc. The interferometric data often reveal small-scale velocity structure, departing from the more gradual large-scale velocity gradients. In some cases, this likely indicates accelerating infall and/or rotational spin-up in the inner envelope; the median velocity gradient from the interferometric data is $\\sim$10.7 \\kmspc. In two systems, we detect high-velocity HCO$^+$ ($J=1\\rightarrow0$) emission inside the highest-velocity \\nthp\\ emission. This enables us to study the infall and rotation close to the disk and estimate the central object masses. The velocity fields observed on large and small-scales are more complex than would be expected from rotation alone, suggesting that complex envelope structure enables other dynamical processes (i.e. infall) to affect the velocity field. ",
        "introduction": "Infall and rotation in dense cores and protostellar envelopes both play important roles in the formation of protostars and their surrounding disks. Infall must be taking place in the envelopes because we observe newborn protostars embedded within their natal clouds. The two classic analytic theories of collapse describe infall as either being outside-in \\citep{larson1969} or inside-out \\citep{shu1977}. The initial angular momentum of the protostellar cloud governs the formation and sizes of the proto-planetary disks \\citep[e.g.][]{cassen1981,tsc1984} and will affect the ability of the cloud core to fragment into multiple stellar systems \\citep[e.g.][]{burkert1993,bonnell1994a}. The ubiquitous formation of proto-planetary disks \\citep{haisch2001,hernandez2007} and the prevalence of binary systems \\citep{raghavan2010} lends strong indirect evidence for the presence of rotation.  Therefore, observing these two processes in protostellar envelopes has enormous potential for constraining star formation theory through comparisons to analytic models and numerical simulations. Dense protostellar cores are well known as sites of isolated, low-mass star formation \\citep[e.g.][]{shu1987,bm1989,mckeeostriker2007} and are the ideal place to study the kinematic structure from large, core/envelope scales ($\\sim$0.1 pc) down to scales near the disk radius.  Several previous studies have attempted to characterize the rotation in dark clouds and dense cores. \\citet{arquilla1986} examined the kinematic structure of dark clouds using CO on $\\sim$10\\arcmin\\ scales, finding that some exhibit possible rotation signatures derived from velocity gradients. Later, \\citet{goodman1993} and \\citet{caselli2002} used the dense molecular tracers \\nht\\ and \\nthp\\ to examine rotation in the dense cores within dark clouds ($\\sim$2-3\\arcmin\\ scales). \\citet{goodman1993} and \\citet{caselli2002} found typical velocity gradients of $\\sim$1 to 2 \\kmspc, which were interpreted to be slow, solid-body core rotation. However, the observations had limited resolution (60\\arcsec-90\\arcsec), nearly the size of the envelopes in some cases. Rotation is not necessarily solid-body; however, this assumption simplifies analysis in terms of equipartition and the early data did not warrant more sophisticated models.  \\citet{caselli2002} noted that their finer resolution as compared to \\citet{bm1989} revealed clear deviations from linearity in the velocity gradients. This indicates that the velocity structure of cores may be more complex than simply axisymmetric, solid-body rotation. \\citet{chen2007} carried out a higher resolution study using interferometric observations of \\nthp\\ from OVRO on a sample of protostars in the Class 0 \\citep{andre1993} and Class I phases \\citep{lada1987}. They found substantially larger velocity gradients in the inner envelopes, with complex velocity fields, only showing probable rotation signatures in a few cases.  While the possible rotation has been probed using optically thin tracers, the study of infall in protostellar envelopes has been limited to observations of optically thick tracers toward the envelope center. \\citet{zhou1993} and \\citet{myers1995} used observations of CS and H$_2$CO to infer the detection of infall in the envelope of the Class 0 protostar B335. The optically thick molecular lines are self absorbed at the line-center and the key signature of infall is the blue-shifted side of the line being brighter than the red-shifted side. Later, \\citet{difrancesco2001} found inverse-P Cygni line profiles in H$_2$CO which are interpreted as infall. Furthermore, infall at large-scales may have been detected in the pre-stellar core L1544 \\citep{tafalla1998}.  These previous studies have interpreted the kinematic structure in terms of axisymmetric envelopes. However, our recent studies \\cite[i.e.][]{tobin2010a,looney2007} have revealed that the dense envelopes surrounding the youngest, generally Class 0, protostars often have complex, non-axisymmetric morphological structure \\citep[also see][]{bm1989,myers1991,stutz2009}. Out of the 22 protostellar systems exhibiting extinction at 8\\mum, only three appeared to be roughly axisymmetric. We suggested that the asymmetric envelope structure may play a role in the formation of binary systems by making fragmentation easier and that the initial disk structure may be perturbed from uneven mass loading. However, in order to understand the effects that complex envelope structure has on disk formation and fragmentation, the kinematics of the dense gas must be characterized.  Given the apparent prevalence of complex morphological structure in protostellar envelopes, the kinematics of the envelopes must also be studied on scales such that the morphological structure is spatially resolved. To examine the kinematic structure of morphologically complex envelopes, we have undertaken a molecular line survey focusing on nearby (d $<$ 500 pc), embedded protostars drawn from our sample of envelopes in \\citet{tobin2010a}. We have approached this study from two directions. First, we obtained new single-dish \\nthp\\ mapping of sixteen systems with the IRAM 30m, a factor of two improvement in resolution over \\citet{caselli2002}. Secondly, we obtained new interferometric \\nthp\\ and \\nht\\ observations of fourteen systems with resolutions between 3.5\\arcsec\\ and 6\\arcsec. The analysis of both the single-dish and interferometric observations offers a comprehensive view of the kinematic and morphological structure from 0.1 pc scales down to $\\sim$1000 AU. Complementary \\textit{Spitzer} imaging gives us a clear view of the outflow angular extent and direction in all of these objects.  In this paper, we are primarily presenting the dataset as a whole and give a basic analysis of each object; an upcoming paper will give a more detailed interpretation of the velocity structures. The data for each object are quite rich and we find many envelopes with velocity gradients normal to the outflow on large-scales in the single-dish maps, persisting down to small-scales in the interferometer maps. We have organized the paper as follows: Section 2 discusses the sample, observations, data reduction, and analysis, Section 3 presents our general results and discusses each object in detail, and Section 4 discusses the basic overall properties of the sample as a whole. ",
        "conclusions": "\\subsection{Large-Scale Velocity Field} The large-scale velocity gradients in dense molecular cores have long been interpreted as rotation (possibly solid-body) \\citep[e.g.][]{goodman1993, caselli2002, belloche2002, belloche2004, chen2007}. If the large-scale motions are indeed rotation, then the circumstellar disk forming from the collapse of the surrounding envelope should be rotating in the same direction as the surrounding cloud, assuming axisymmetric collapse \\citep[e.g.][]{shu1987,bodenheimer1995}. However, if the core has fragmented, then the angular momenta of the individual disks could be misaligned, but the sum of the angular momenta would still be aligned with the cloud rotation. The direction of the jet and outflow of the system is commonly thought to reflect the current angular momentum vector of the protostar and disk system. This is because outflows and jets are thought to be related to the magnetic field of the protostar and disk which should be along the rotation axis \\citep{pudritz1983,shu1994} and resolved observations of disks have shown that jets are oriented normal to the disk midplane. Thus, we have used the outflow direction and protostar locations as guideposts as to where we should measure the velocity gradients relative to and in which direction because the equatorial plane of the envelope is where we are most likely to observe rotation.  The 2D fitting method (see Section 2.6.2) shows that eleven out of sixteen systems have a gradient direction that is within 45\\degr\\ of normal to the outflow axis in the single-dish data; however, there is considerable spread in this distribution. We caution that the PA of the gradients calculated by this method may have systematic error due to the velocity fields not being uniform and the velocities are sampled over a region that is not symmetric, leading to a bias of data points in a particular direction. Visual inspection of the velocity fields does yield a similar result to the 2D fitting, with eleven normal to the outflow and five not normal; IRAS 03282+3035, HH270 VLA1, L1527, L1521F, and L483 were not normal to the outflow. In contrast to the 2D fitting, we visually identify L483 and HH270 VLA1  as not normal to the outflow and L1157 and Serpens MMS3 to be normal. We further caution that most velocity fields have structure not easily described by a single position angle or velocity gradient. Nevertheless, as a simple significance test of the gradient-outflow direction relation, we can apply the binomial distribution if we consider the $<$45\\degr\\ and $\\ge$45\\degr\\ as two bins and that the gradient direction relative to the outflow may be oriented randomly between 0 and 90\\degr\\ with a mean of 45\\degr\\ \\citep{bevington1969}. Thus, the probability of the relative directions falling within $<$45\\degr\\ or $>$45\\degr\\ should be equal. The chance of eleven or more objects out of sixteen with randomly oriented velocity gradients falling within the $\\ge$45\\degr\\ bin is only $\\sim$10\\%.  We also attempted to see if there is any trend with envelope mass and velocity gradient, we compared the observed velocity gradients and linewidths with the escape velocity in Figure \\ref{bound}. The envelope mass is calculated from the dust mass measured from 8\\mum\\ extinction (Paper I, Table 1) plus 1 $M_{\\sun}$ to account for the central object; note that the central object masses may be overestimated while the envelope masses are likely underestimated (Paper I). There is no apparent trend of velocity gradients or linewidth with increasing envelope mass (in terms of escape velocity). This figure further demonstrates that these protostellar systems are consistent with being gravitationally bound and are not supported by rotation or turbulence, consistent with many previous studies \\citep[e.g.][]{goodman1993,caselli2002,chen2007}.  Furthermore, the smallest velocity gradients in our sample appear to be coming from the most symmetric envelopes L1521F, IRAS 16253-2429, and L1157, and the largest gradients are found in the morphologically complex HH211 and IRAS 03282+3035 systems. This could be taken to mean that the velocity structure is no strongly influenced by complex projection effects in the symmetric systems and are observing slow envelope rotation, while the more complex systems have projection effects altering the observed velocity structure. However, despite this trend, we note that the envelope of the low-luminosity source IRAM 04191 is approximately symmetric (Paper I) and has a velocity gradient of 17 \\kmspc\\ \\citep{belloche2002, belloche2004}, a clear counter-example to this morphological trend.  The simple interpretation of our velocity gradient data as a whole is that we are observing core rotation in these systems and that they have a variety of angular momenta ; some angular momenta being quite large (Figure \\ref{sdgradienthistos}). However, if we were observing pure rotation on large-scales, then one would expect the velocity gradient directions to be much more clustered toward being orthogonal to the outflow rather than the broad distribution shown in Figure \\ref{paoffset}. Therefore, either the complex morphology makes the rotation ambiguous, causing the velocity gradients to not be orthogonal to the outflow or we are not observing pure rotation. In either case, the angular momenta derived from the velocity gradients will be suspect at best. Given these complicating factors, we regard the velocity gradients taken as 1D cuts normal to the outflow as most likely to be probing the velocity gradients due to rotation, but these values should be regarded as upper limits.  If we are not observing pure rotation, the only other dynamical process which should give rise to ordered velocity structures are infall and/or outflow entrainment. The outflows from the protostars in our sample are highly collimated, limiting their ability to affect the large-scale kinematic structure, see Section 4.3 for further discussion, which leaves infall as the only other mechanism to contribute to the velocity field. This would require that infall is happening on large-scales and that collapse would need to be outside-in and not inside-out. Large-scale infall is shown to be possible in numerical simulations of complex, filamentary cores forming within a molecular cloud by \\citet{smith2011}. Their simulations show that there is infall onto filaments (scales of 0.1 pc to 0.01 pc) from the surrounding molecular cloud, with subsequent infall from the filament to the protostar (sink particle) ($<$ 0.01 pc). Thus, infall convolved with rotation in real envelopes could lead to some of the complex velocity structures that we observe.  In summary, it is difficult to interpret the large-scale velocity gradients wholly being due to rotation since the envelopes are asymmetric and the gradients are often not normal to the outflow. Furthermore, flows or infall along a filament could give a similar signature to rotation simply from geometric projection, and if an envelope is filamentary, infall along the envelope would also produce a velocity gradient normal to the outflow. We therefore suggest that a component of infall velocity could be projected along our line of sight, entangled with the rotation velocity, resulting in the large velocity gradients. This issue of rotation versus projected infall will be further explored in an upcoming paper.   \\subsection{Small-scale Velocity Structure}  The high-resolution interferometer data are essential for probing the kinematics of the envelope at scales smaller than 5000 AU. These data enable us to localize the gas in the envelope and to assess the dynamical processes at work. Small-scale velocity sub-structure, beyond an ordered/linear velocity gradient, is found in the interferometric observations of most envelopes, including L1157, L1165, HH108IRS, Serpens MMS3, L1527, CB230, IRAS 03282+3035, and RNO43. These features are sometimes apparent in the velocity maps or profiles shown in Figure \\ref{profiles}, but they also appear as increased linewidth in the inner envelope. The small-scale velocity structure is generally found on $\\sim$2000 AU scales in the envelopes. This radius could possibly be the centrifugal radius where material can be rotationally supported against gravity and the increased rotation velocity makes this region stand out against the rest of the envelope in velocity. However, assuming the large-scale velocity gradients reflect rotation, there would not be enough angular momentum for material to be rotationally supported at this radius.  We found that the relationship between outflow axis and velocity gradient direction in the interferometer data trends even more strongly toward being normal to the outflow that in the single-dish data. Twelve of fourteen systems have gradients within 45\\degr\\ of the outflow in the right panel Figure \\ref{paoffset}, as found by the 2D fitting technique. Using the same statistical analysis as the single-dish gradients, this result being due to chance is $\\sim$0.6\\%. Note that this plot is missing L1157 and Serpens MMS3 since their complex velocity fields could not be reliably fit. Visually, we find that twelve envelopes clearly have gradients normal to the outflow and four do not (HH108MMS, L483, IRAS03282, and HH270VLA1); L483 and HH270 VLA1 were counted as within 45\\degr\\ of normal to the outflow. Note that RNO43, Perseus 5, and L1521F appear to have gradients normal to the outflow, but they also have complex velocity fields. L1157 is left out of this analysis because the gradient directions from the CARMA and VLA data differ by $\\sim$80\\degr.  The velocity gradient direction is generally consistent between both large and small scales in the single-dish and interferometer data as shown in Figure \\ref{sdintoffset}. This could mean that the dynamical processes observed at large-scales are also responsible for the kinematics observed at small-scales. The protostellar systems which show substantial deviation from large to small-scales are L1527 and L1521F. The difference in L1527 marks a velocity gradient reversal from large to small-scale (as also shown in Figure \\ref{profiles}) and the difference in L1521F reflects that the kinematic structure was not well-resolved in the single-dish data. The \\nthp\\ and \\nht\\ lines in L1521F are also extremely optically thick making the velocity field derived from fitting uncertain since we cannot probe all the gas along the line of sight.  The preference for the vast majority of velocity gradients to be within 45\\degr\\ of normal to the outflow contrasts with \\citet{chen2007}, where only two of nine targets had this feature. Moreover, in \\citet{volgenau2006}, only one protostellar core of three observed had a well-ordered velocity structure. This result may be due to environment since the ordered structure was found in L1448 IRS3, a more isolated system, and the other sources were located in the more complex environment of NGC1333. Our sources in contrast are generally isolated, like those in \\citet{chen2007}. As noted in the preceding paragraphs, we do see several objects with similarly complex velocity fields in our data; however, our greater number of observed systems likely enabled us to find more with ordered velocity fields.  The velocity gradients fit for the interferometer data are systematically larger than those found for the single-dish data, with the average being 8.6 \\kmspc\\ from 1D fitting, the distribution is shown in Figure \\ref{intgradienthistos}. This factor is nearly equal to our increase in resolution in the interferometer data as compared to the IRAM 30m, but the object-to-object increase is more varied. Some of the small-scale velocity structure is due to outflow interactions in the inner envelope as we will discuss in more detail in the following section; however, we have attempted to mask out these regions when fitting the velocity gradients.  Like the single-dish data, we regard the 1D fitting method to more accurately reflect the velocity gradient intrinsic to the envelope itself since the 2D method will be more susceptible to outflow effects on the envelope kinematics. The only objects not showing large increase in the velocity gradient on small-scales are HH211 and IRAS 16253-2429 whose small-scale velocity gradient may be slightly overestimated.  We strongly cautioned in the previous section about interpreting the velocity gradients as rotation on large-scales and we are again hesitant to interpret the small-scale gradients as rotation and/or spin-up. This is because most envelopes are highly filamentary and on small-scales projections effects on the velocity structure will be even more apparent since both infall and rotation velocities increase at small radii. The velocity fields themselves on small-scales are not well-ordered as one might expect from rotationally dominated motion. Highlighting a few examples from Section 3.2: Serpens MMS3 shows deep red-shifted emission and increased linewidth only on one side of the protostar, the velocity gradient in HH108IRS reverses itself just past the protostar, the gradient in L1527 on small-scales is opposite of large scales, and RNO43 has an abrupt velocity jump across the envelope. The complex nature in the velocity fields of many sources do not necessitate an interpretation as rotation. Furthermore, infall velocities will always be necessarily be larger than rotation since the envelopes are not rotationally supported \\citep{chen2007}. Thus, the only way to robustly separate infall velocities from rotation is at $<$ 1000 AU scales, near the centrifugal radius where material can be rotationally supported.  Two systems in our sample (L1165 and RNO43) show high-velocity line wings in the inner envelope that could indeed reflect significant rotation. In both sources, the emission appears to be coming from a radius of $\\sim$600 AU. Our data indicate that the observed velocities could reflect rotationally supported motion around $\\sim$0.5$M_{\\sun}$ central objects. Higher resolution and higher signal-to-noise data are needed to accurately centroid the high velocity emission in order to more precisely constrain the enclosed masses. In the future, ALMA could be used to spatially resolve the high-velocity emission in order to fully separate rotation from infall. Further analysis and modeling of the velocity structures observed with the interferometric data with the goal of determining the kinematic processes at work in the envelopes will be presented in an upcoming paper (Tobin et al. 2011 in preparation).  \\subsection{Outflow Induced Kinematic Structure}  Our knowledge of protostellar outflows has been greatly enhanced in recent years \\citep[e.g.][and references therein]{bachiller1996,arce2007} and their possible effects on the surrounding envelope have been characterized \\citep{arce2006}. Furthermore, IRAC imaging from the \\textit{Spitzer Space Telescope}, in addition to near-IR imaging from the ground, can give a strong constraint on the outflow axis and cavity width \\citep[e.g.][]{seale2008}. These data complement observations of outflow tracers such as CO and give a more complete picture of how the outflow may be impacting the protostellar envelope.  Several protostars in the sample have velocity structures that are strongly suggestive of outflow effects on the \\nthp\\ and \\nht\\ gas kinematics.  L1157, HH108MMS, and Perseus 5 appear to be significantly affected by the outflow, while HH108IRS, HH211, IRAS 03282+3035, L1152, and IRAS 04325+2402 only appear to be mildly affected. In the mildly affected cases, the large-scale bulk motion of the envelopes (line-center velocity) does not appear to be affected by the outflow, rather we see the effects in the linewidth at both large-scales and small-scales. This probably means that only a small portion of the total cloud mass is being affected by the outflow. The bulk motion effects are revealed in the strongly affected cases, generally on small-scales and only visible in the interferometer data.  \\citet{arce2006} presented an empirical model for how the outflow will affect the envelope during protostellar evolution, concluding that the outflow is ultimately responsible for disrupting the protostellar envelopes. The number of objects we find showing outflow effects in the kinematic data strongly support the outflow-envelope interaction framework put forward by \\citet{arce2006} and enable us to offer some further input to this empirical framework.  Perseus 5 and HH108MMS appear to be some of the youngest objects in our sample, as evidenced by their deeply embedded nature and lack of visible outflow cavities in 3.6\\mum\\ or Ks-band imaging. In these systems, the outflow seems to be having the greatest impact on the kinematics in both line-center velocity and linewidth. Therefore, the effects of the outflow on the kinematic structure may be most prominent during its initial breakout of the envelope, early in the Class 0 phase. Smaller-scale effects on the envelope can clearly be seen in the case of L1157 where the envelope material may be entrained at 1000 AU scales. Thus, the outflows could be carrying significant momentum at wide angles near the protostar in order to be actively forcing inner envelope material out. If the outflows are impacting the envelope at small-scales and wide angles, an important question remains as to how much the outflow can quench infall in filamentary envelopes. In L1157, the amount of entrained gas appears tenuous, only $\\sim$0.04 $M_{\\sun}$ out of 0.7 $M_{\\sun}$ in the inner envelope, following a simple analysis \\citep{goldsmith1999} assuming an \\nthp\\ abundance of 10$^{-9}$ per H$_2$. On larger-scales ($>$ 1000AU), we expect that material extended normal to the outflow in filamentary envelopes will not be strongly influenced by the outflow, as the outflows appear to stay well collimated throughout the Class 0 phase.  \\subsection{Linewidths}  The linewidths in the envelopes are generally quite small away from the protostar. In most cases, the linewidths are 0.2 - 0.5 \\kms\\ in the single-dish data; linewidths averaged over the entire source are given Table 3. Note that these linewidths are substantially broader than the 10K thermal linewidth of \\nthp\\ which is $\\sim$0.13 km/s, and is generally attributed to turbulent motions and/or unresolved velocity gradients along the line of sight. However, these linewidths are not large enough to make the envelope unbounded, as shown Figure \\ref{bound}.  The interferometer data find even smaller linewidths toward the outer edges of the envelopes; there were several cases where the lines are less than 2 channels wide. The narrow line-widths in the interferometer data likely stem from the larger scale structure being resolved out; the larger scale emission may have some turbulent or large-scale infall velocity component. The regions of broad linewidth in the interferometer observations appear to be directly related to the outflow, increased line-of-sight motion, and/or heating from the protostar; not a transition to a turbulent core \\citep[cf.][]{chen2007}.  \\nht\\ linewidths were used by \\citet{pineda2010} to probe the kinematics of the large-scale molecular core, detecting a transition from the quiescent core to the turbulent cloud. We looked for such an effect in our sensitive \\nthp\\ data but did not find similar structure. We also did not detect \\nthp\\ emission on the scales for which \\citet{pineda2010} were able to detect \\nht, despite our high sensitivity. We believe that this is likely due to the differing critical densities between \\nthp\\ and \\nht\\ ($\\sim$2$\\times$10$^3$ cm$^{-3}$ for \\nht\\ (1,1) and $\\sim$1.4$\\times$10$^5$ cm$^{-3}$ for \\nthp\\ ($J=1\\rightarrow0$)).  We did, however, find that some maps (i.e. L673, HH211, HH108, RNO43) had regions with two distinct velocity components. The regions of overlap between the components appear as artificially large linewidths in the maps generated by the hyperfine fitting routine (section 2.6.1). We show three \\nthp\\ spectra from HH211 taken at three positions showing the different velocity components in Figure \\ref{lineprofiles}; this is indicative of what takes place in the other regions showing this feature. Notably, the additional velocity components tend to appear toward the edge of the maps, except in RNO43 where the transition takes place near the protostar. The distinct velocity components are located in regions that appear form a contiguous structure when viewed in 8\\mum\\ extinction. The second velocity component generally appears about $\\sim$0.05 pc from the nearest protostar. Thus, the reason for multiple components could have to do with the initial conditions of the clouds themselves. Recent simulations have suggested that colliding clouds could be an important component to setting up the initial conditions for star formation \\citep[e.g.][]{heitsch2006}.   \\subsection{Chemical Effects on Molecular Tracers}  The kinematic data presented are based on the molecular tracers \\nthp, \\nht, and \\hcop. We have used \\nthp\\ and \\nht\\ relatively interchangeably, since they appear to trace the same kinematics and physical conditions (see Section 3.1 and \\citet{johnstone2010}). This makes sense because we know that in the pre-stellar phase, the formation of \\nthp\\ and \\nht\\ appear to be linked given their similar abundance distributions and only deplete onto dust grains at very high densities \\citep{bergin2007}. In many of our observations, \\nthp\\ and \\nht\\ only appear to trace the gas on scales $>$1000 AU from the protostar; the emission generally peaks near the protostar, but not directly on it. This indicates a drop in abundance either due to depletion or destruction of the molecules. Several of the protostars for which there are both interferometric \\nht\\ and \\nthp\\ observations (in this paper or in the literature) show decreased emission at the location of the protostar in both \\nthp\\ and \\nht. Observations of CB230 and IRAS 03282+3035 in \\nht\\ (Figures \\ref{CB230} and \\ref{IRAS03282}) and \\nthp\\ \\citep{chen2007} are clear examples of \\nht\\ and \\nthp\\ not being peaked coincident with the protostar. L1157 also exhibits this effect in \\nthp\\ from the PdBI observations and also does in \\nht\\ if the longer baselines are given more weight in the VLA map. Finally, the low-luminosity source IRAM 04191 shows a similar depletion pattern in both of these tracers \\citep[][J. Mangum, Private Communication]{belloche2004}.  We can understand the decrease in \\nthp\\ emission in terms of molecular destruction by reactions with other molecules. In the inner envelope, where the temperatures rise above 20K, the CO that has depleted onto the dust grains \\citep{bergin2002} is released back into the gas phase. CO and \\nthp\\ rapidly react to form \\hcop; this is the dominant destruction mechanism for \\nthp\\ \\citep{aikawa2001,lee2004}. This is the same reason that \\nthp\\ is not seen in outflows and why the \\nthp\\ is often not centrally peaked on the protostars in our interferometric observations. \\nht\\, however, is not directly destroyed by CO, but rather \\hcop\\ \\citep{lee2004}. \\hcop\\ will readily form within the region of CO evaporation making it available to react with \\nht. Alternatively, \\nht\\ could also become depleted onto dust grains in an ice mantle, and if it is well-mixed with the water ice, then it would only be evaporated at temperatures $\\geq$100K. \\nht\\ ice is frequently observed in the envelopes surrounding protostars \\citep{bottinelli2010} via mid-infrared spectroscopy. The absorbing \\nht\\ ice should be in the inner envelope since \\nht\\ only depletes onto grains at high densities. If any of the \\nht\\ ice were released into the gas phase, then \\hcop\\ would be present to destroy it.  Since both \\nht\\ and \\nthp\\ are not present on scales $<$1000 AU, we must look for other tracers to probe the kinematic structure on these scales. In two protostars, we have been able to use \\hcop\\ to trace the small-scale kinematic structure. The \\hcop\\ emission from L1165 and RNO43 show high-velocity wings on small-scales inside the innermost \\nthp\\ emission; the centroids of the red and blue-shifted emission are offset from the protostar, normal to the outflow direction, at radii of $\\sim$600AU.  Chemical models of an infalling protostellar envelope have been calculated by \\citet{lee2004}, showing that \\hcop\\ becomes enhanced at small-scales after formation of the protostar. This reflects the evaporation of CO ice, releasing CO back into the gas phase where \\hcop\\ is readily formed. In addition, the primary destruction pathway of \\nthp\\ from CO results in the formation of HCO$^+$ and N$_2$. Thus, we can see why the \\nthp\\ is tracing the gas at larger radii and lower velocities, while in the inner envelope HCO$^+$ is readily being formed and can then trace the small-scale high-velocity gas.  Our dataset explicitly shows how multiple tracers can be used to gain a more complete picture of the kinematics in protostellar envelopes. \\nthp\\ and \\nht\\ are excellent tracers of the cold, dense gas on scales from $\\sim$1000-2000 AU out to $\\sim$10000-20000 AU. Inside of 1000AU, an abundant tracer of the warm, inner envelope tracer must be observed; HCO$^+$ works quite well in two out of nine sources observed with CARMA and other tracers and/or higher-J transitions of \\hcop\\ may work as well (see \\citet{lee2009} and \\citet{brinch2007}). However, these other tracers are often found in outflows and they must be observed with sufficiently high resolution to confirm their origin in the envelope and not the outflow. In the future, ALMA may be able to observe inner envelope tracers to resolve the motion of the dense gas in the inner envelope, tracing infall onto the disk."
    },
    "1107/1107.5771_arXiv.txt": {
        "abstract": "The physical nature of compressible turbulence is of fundamental importance in a variety of astrophysical settings. We present the first direct evidence that mean kinetic energy cascades conservatively beyond a transitional ``conversion'' scale-range despite not being an invariant of the compressible flow dynamics. We use high-resolution three-dimensional simulations of compressible hydrodynamic turbulence on $512^3$ and $1024^3$ grids. We probe regimes of forced steady-state isothermal flows and of unforced decaying ideal gas flows. The key quantity we measure is pressure dilatation cospectrum, $E^{PD}(k)$, where we provide the first numerical evidence that it decays at a rate faster than $k^{-1}$ as a function of wavenumber. This is sufficient to imply that mean pressure dilatation acts primarily at large-scales and that kinetic and internal energy budgets statistically decouple beyond a transitional scale-range. Our results suggest that an extension of Kolmogorov's inertial-range theory to compressible turbulence is possible. ",
        "introduction": "Turbulence plays a critical role in essentially all astrophysical systems that involve gas dynamics. For example, interstellar turbulence is believed to be driven on large scales by differential galactic rotation or supernovae explosions and dissipated at the smallest scales by microphysical processes. It is widely believed that energy is transferred from the largest scales down to the dissipation scales through a cascade process. Measurements of interstellar scintillation (ISS) (\\cite{Armstrongetal81}) caused by scattering in the interstellar medium shows a power-law electron density spectrum with a slope close to Kolmogorov's $-5/3$ scaling over 5 decades in scale. A composite spectrum combining ISS data with that of differential Faraday rotation angle and gradients in the average electron density also yields a Kolmogorov power-law over at least ten decades in scale, dubbed as ``the great power-law in the sky'' (\\cite{Armstrongetal95}). High-resolution simulations of the interstellar medium (\\cite{Kritsuketal07}) also reveal power-law scaling of spectra and structure functions which are reminiscent of incompressible turbulence (albeit with different slopes). Yet, observations of the interstellar medium suggest that the Mach number of turbulent motions is of order $0.1$ to $10$, such that the flow is often compressible (\\cite{ElmegreenScalo04}). Similarly, high-resolution simulations of the intracluster medium (ICM) between galaxies (\\cite{Xuetal09,Xuetal10}) reveal Kolmogorov-like spectra of kinetic energy at subsonic and transonic Mach numbers. It is not clear how Kolmogorov's 1941 phenomenology, which forms the cornerstone for our understanding of incompressible hydrodynamic turbulence, can carry over to such flows which exhibit significant compressibility effects. Several recent studies have addressed the problem of compressible hydrodynamic turbulence numerically (e.g. \\cite{Porteretal98,Kritsuketal07,Schmidtetal08,Federrathetal10,PetersenLivescu10}) and theoretically (e.g. \\cite{RubinsteinErlebacher97,Boldyrev02,Falkovichetal10}), but a physical description equivalent to that of Kolmogorov's still eludes us.  The idea of a cascade itself is without physical basis since kinetic energy is not a global invariant of the inviscid dynamics. Compressible flows allow for an exchange between kinetic and internal energy through two mechanisms: viscous dissipation and pressure dilatation. While the former process is localized to the smallest scales just like in incompressible turbulence, the latter is a hallmark of compressibility and can {\\it a priori} allow for an exchange at any scale through compression and rarefaction. Recently, \\cite{Aluie11c,Aluie11b} showed how the assumption that pressure dilatation cospectrum decays fast enough rigorously implies that mean kinetic energy cascades conservatively despite not being an invariant. The assumption entails that the mechanism of pressure dilatation acts primarily at the largest scales and vanishes \\emph{on average} at smaller scales. In this Letter, we shall provide the first empirical evidence in support of this assumption.  The outline of this Letter is as follows. Section \\ref{Numerics} describes our numerical simulations and Section \\ref{Coarse-graining} presents our coarse-graining method for analyzing nonlinear scale interactions. Section \\ref{sec:PressDil} discusses the significance of pressure dilatation cospectrum and the assumption we are testing. Our main results are then described in Section \\ref{sec:Results} followed by a discussion on the role of shocks in Section \\ref{sec:Shocks}. The Letter concludes with Section \\ref{sec:Conclusion}. ",
        "conclusions": "\\label{sec:Conclusion} The main result of our Letter was to provide the first empirical evidence that pressure dilatation co-spectrum decays at a rate faster than $k^{-1}$, in accord with condition (\\ref{scaling4}). This is sufficient to imply that exchange of kinetic and internal energy through compression and rarefaction takes place at the large scales, on average. At smaller scales beyond a transitional ``conversion'' scale-range, mean kinetic and internal energy budgets statistically decouple and kinetic energy can only reach dissipation scales via a scale-local conservative cascade process. The existence of such a conservative cascade in compressible turbulence justifies expectations that spectra with power-law scalings should exist in such flows. Furthermore, it suggests that  observations of power-law spectra in astrophysical systems are indeed evidence of a turbulent cascade process similar (although not necessarily identical) to that in incompressible flows.  This work leads us to some new questions which we hope to address in future studies. Does the decay rate of pressure dilation cospectrum depend on the decaying/forced nature of the turbulence or on other parameters of the flow such as equation of state, Mach number, and the ratio of compressive-to-solenoidal components of the velocity field? Is the power law scaling of pressure dilation cospectrum that we find in Figure \\ref{Fig:isoeosSpecPD} reflecting the asymptotic scaling at arbitrarily high Reynolds numbers? Does the value of $PD(\\ell)$ in Figure \\ref{Fig:isoeosSGSFlux} relative to that of SGS flux increase by increasing the compressive amplitude of the forcing? We invite future studies to investigate these questions and to try and reproduce our results presented in this Letter. Verifying assumption (\\ref{scaling4}) under a variety of controlled conditions would substantiate the idea of statistical decoupling between kinetic and internal energy budgets. This is of paramount importance for future attempts to extend Kolmogorov's ideas on a conservative cascade of kinetic energy to compressible turbulence.  \\vspace{0.4cm} \\noindent {\\small {\\bf Acknowledgements.} We are grateful to J. Cho for providing us with his forcing subroutine. H. A. thanks R. E. Ecke, G. L. Eyink, S. S. Girimaji, S. Kurien, and D. Livescu for useful discussions. H.A. acknowledges partial support from NSF grant PHY-0903872 during a visit to the Kavli Institute for Theoretical Physics. This research was performed under the auspices of the U.S. DOE at LANL under Contract No. DE-AC52-06NA25396 and supported by the LANL/LDRD program and by the DOE/Office of Fusion Energy Science through NSF Center for Magnetic Self-Organization.}"
    },
    "1107/1107.2198_arXiv.txt": {
        "abstract": "High sensitivity observations of radio halos in galaxy clusters at frequencies $\\nu\\le 330$ MHz are still relatively rare, and very little is known compared to the classical 1.4 GHz images. The few radio halos imaged down to 150--240 MHz show a considerable spread in size, morphology and spectral properties. \\\\ All clusters belonging to the GMRT Radio Halo Survey with detected or candidate cluster--scale diffuse emission have been imaged at 325 MHz with the GMRT. Few of them were also observed with the GMRT at 240 MHz and 150 MHz. For A\\,1682, imaging is particularly challenging due to the presence of strong and extended radio galaxies at the center. Our data analysis suggests that the radio galaxies are superposed to very low surface brightness radio emission extended on the cluster scale, which we present here. ",
        "introduction": "The origin of radio halos and relics is strictly connected to the dynamics of the hosting cluster. Recent works, based on the GMRT Radio Halo Survey (Venturi et al. 2007 \\& 2008), have highlighted a number of relevant statistical properties of radio halos, such as the {\\it bimodality}: clusters are either {\\it radio loud}, in which case the halo radio power correlates with the cluster X--ray luminosity, following the well--known logL$_{\\rm X}$--logP$_{\\rm 1.4~GHz}$ correlation, or {\\it radio quiet}, with radio power upper limits well below the same correlation  (Brunetti et al. 2009). Moreover, a quantitative radio/X--ray analysis of the radio halo/cluster merger connection shows that halos are found only in cluster mergers; relaxed clusters never host them, and only few massive, luminous and merging clusters are void of diffuse emission at the sensitivity level of the current radio interferometers (Cassano et al. 2010)  These results support  the idea that radio halos are the result of relativistic particle re--acceleration by turbulence induced in the cluster volume by massive cluster mergers (the {\\it re--acceleration model}, Petrosian 2001 and Brunetti et al. 2001). This model further predicts the existence of a population of halos with ``ultra--steep spectrum'', i.e. $\\alpha \\ge 1.5$, which would be connected to the  less energetic, but much more frequent minor mergers (Cassano 2009 and Cassano, these proceedings). A\\,521, with a spectral slope $\\alpha \\sim 1.9$, is the first ultra--steep spectrum radio halo (USSRH) discovered (Brunetti et al. 2008, Dallacasa et al. 2009). Due to their spectral properties, USSRH are best detectable at low frequencies, i.e. $\\nu\\le$330 MHz.  To improve the knowledge of the low--frequency properties of radio halos, and possibly identify candidate USSRHs, all clusters in the GMRT radio halo survey with detected diffuse or candidate cluster--scale emission were observed with the GMRT at 325 MHz, and some of them were followed up at 240 and 150 MHz. Some results have already been published in recent papers (Giacintucci et al. 2009 \\& 2011, Macario et al. 2010 \\& 2011, Venturi et al. 2011). Here we present and briefly discuss the challenging case of A\\,1682. ",
        "conclusions": ""
    },
    "1107/1107.0965_arXiv.txt": {
        "abstract": "Radio galaxies with bent jets are predominantly located in groups and clusters of galaxies.  We use bent-double radio sources, under the assumption that their jets are bent by ram-pressure, to probe intragroup medium gas densities in galaxy groups.  This method provides a direct measurement of the intergalactic gas density and allows us to probe intergalactic gas at large radii and in systems whose intragroup medium (IGM) is too cool to be detected by the current generation of X-ray telescopes.  We find gas with densities of $10^{-3}-10^{-4}\\cmc$ at group radii from $15-700$ kpc.  A rough estimate of the total baryonic mass in intergalactic gas is consistent with the missing baryons being located in the intragroup medium of galaxy groups.  The neutral gas will be easily stripped from dwarf galaxies with total masses of $10^{6}-10^{7}\\ \\msol$ in the groups studied here.  Indications are that intragroup gas densities in less-massive systems like the Local Group should be high enough to strip gas from dwarfs like Leo T and, in combination with tides, produce dwarf spheroidals. ",
        "introduction": "Galaxy groups are ubiquitous, intermediate density structures, spanning the range between isolated field galaxies and rich clusters \\citep{1983ApJS...52...61G,1987ApJ...321..280T,2004MNRAS.348..866E,2010A&A...514A.102T}.   According to the hierarchical scenario of the formation of large-scale structure, groups are the building blocks of rich clusters of galaxies \\citep{2008A&A...489...11B,2009MNRAS.400..937M}.  Thus, understanding the physical properties of clusters and their member galaxies requires knowledge of the extent to which galaxy evolution occurs in the group environment.  The effectiveness of physical processes that may be responsible for this evolution, like ram-pressure stripping and strangulation, is a function of the density of intergalactic gas.  There are only a handful of ways to detect this gas, namely X-ray observations which are limited to the more massive groups and UV absorption line studies which probe gas with a limited temperature range.  We discuss here a complementary method, applicable in even low mass groups, which can determine the intergalactic gas density using bent-double radio sources.  Groups are likely to contain a significant fraction of the baryons in the local universe \\citep{1998ApJ...503..518F,2000ARA&A..38..289M}.  Attempts to reconcile the baryon content of the local universe with the baryon density seen at high redshift have failed to locate the majority of baryons \\citep{2004ApJ...616..643F}.  Specifically, the baryon deficit appears to scale with potential well depth such that the baryon content of massive clusters is nearly as expected while groups and individual galaxies are lacking \\citep{2003ApJ...585L.117B, 2010ApJ...719..119D}.  These locally ``missing baryons'' are predicted by simulations to exist in a warm-hot intergalactic medium (WHIM) that pervades large-scale structures \\citep{1999ApJ...514....1C,2001ApJ...552..473D}.  In simulations this gas is shock heated during structure formation, leading to temperatures in the range $10^5 < T < 10^7$ K and a broad density distribution that is peaked at $10-20$ times the critical density of the universe \\citep{2006ApJ...650..560C}.  Ultraviolet (UV) and X-ray absorption line detections certainly confirm the existence of the highly ionized WHIM (see the review by \\citet{2007ARA&A..45..221B}) but the spatial distribution of this gas inside and outside of galaxies is still unconstrained.  We explore the potential impact of the presence of a widespread intergalactic medium with with densities $\\sim 10^{-3}-10^{-4} \\cmc$ by looking at the fate of the gas content of dwarf spheroidal galaxies.  Dwarf spheroidal (dSph) galaxies are dark-matter dominated, gas poor, low surface brightness, and have stellar populations with old to intermediate ages.  Evidence for the role of environment in the evolution of dSphs is seen in the morphology-density relations which exist for the Local Group and nearby groups \\citep{1994AJ....107.1328V,2009AJ....138.1037C,2003AJ....125..593S,2009AJ....137.3038B}; dSphs are preferentially found near large galaxies.  Their progenitors are thought to be similar to dwarf irregular (dIrr) galaxies which were processed by their environment, causing angular moment loss and gas stripping \\citep{2003AJ....125.1926G,2010AdAst2010E..25M}.  This is supported by the observation that structurally, when comparing surface brightness profile shapes, dSphs, dIrrs, and galaxy disks form a family that is distinct from that of elliptical and dwarf elliptical galaxies \\citep{1983ApJ...266L..21L,1985ApJ...295...73K,1987nngp.proc..163K,2009ApJS..182..216K}.  However, dSphs have slightly higher metallicities and surface brightnesses than current dIrrs, indicating that they have undergone comparatively more early and rapid star formation episodes \\citep{1998ARA&A..36..435M}.  A third class of dwarf, known as ``transition-type'', possesses a mixture of dIrr and dSph properties and is likely to be the progenitor of future dSph galaxies.  Interestingly, these transition dwarfs also show a morphology-density relation in the Centaurus A group, with average distances from large galaxies intermediate between those of dSphs and dIrrs \\citep{2009AJ....138.1037C}.  Ram pressure stripping, which depends on the density of the intergalactic gas in these groups, is the favored mechanism for the removal of the interstellar neutral gas from dSph progenitors \\citep{2003AJ....125.1926G,2010AdAst2010E..25M}.    Radio galaxies with bent jets are preferentially located in groups and clusters of galaxies, whereas straight and single component radio sources are not \\citep{2011AJ....141...88W}.  We present measurements of intergalactic gas densities in groups of galaxies, assuming that the jets of these radio galaxies are bent backward by ram-pressure, for five new sources.  These are groups for which no X-ray data exist or which are undetected in current observations.  We include the two radio galaxies from \\citet{2008ApJ...685..858F} (hereafter F08) as we have slightly altered our original set of assumptions about the bulk flow speeds along the radio galaxy jets.  The original results change only marginally whereas the derived densities are now more constrained.  Sources have been renamed, as follows, to allow easy referencing:   CGCG 156-060 (S3), SDSS J154849.35+361035.3 (S4), SDSS J112038.52+291234.1 (S5), SDSS J115434.81+363539.8 (S6), SDSS J212616.07-071046.3 (S7) and FIRST J124942.2+303838 (S1), SDSS J085331.86+23400.0 (S2) from F08.   Table \\ref{tab:sour} includes source coordinates and properties as well as results.  We use a Hubble constant of $73 \\kmsmpc$ and velocities which are corrected with respect to the Cosmic Microwave Background when determining distances.  This Hubble constant corresponds to $\\rho_{crit}=3c^2H_0^2/8\\pi G=1.0 \\times 10^{-29} \\gcmc$. ",
        "conclusions": "In the previous section, and Table \\ref{tab:sour}, we have presented measurements of intergalactic gas densities using radio sources located in groups of galaxies.  Assuming that these densities are representative of the IGM in groups in general we now explore the impact this gas can have on gas rich dwarf galaxies, hot gas in galaxy halos, and the baryon fraction.  \\subsection{Ram Pressure Stripping of Dwarf Galaxies}  There are a number of indications that the group environment has an affect on dwarf galaxies.  In combination, dSph galaxies in the Local, Centaurus A, and Sculptor groups have mean distances of $0.23 \\pm 0.20$ Mpc from the nearest large spiral galaxy, while transition-type and dIrr galaxies have mean distances of $0.54 \\pm 0.31$ Mpc and $0.85 \\pm 0.55$ Mpc, respectively \\citep{2009AJ....138.1037C}.  These transition dwarfs are basically gas-rich dSphs and accordingly they fall on the luminosity-metallicity diagram near other dSphs \\citep{1998ARA&A..36..435M}.  \\citet{2009ApJ...696..385G} examine the \\hi content of Local Group dwarf galaxies and find that within $\\sim 270$ kpc, of either the Milky Way or Andromeda, all dwarfs (except more massive dwarf ellipticals) are undetected in \\hi while most galaxies at larger radii have $10^5 - 10^8\\ \\msol$ of neutral gas.  The flat (compared to the field \\himf) low mass slope of the \\himf for the group environment \\citep{2009MNRAS.400.1962K,2009AJ....138..295F, 2005ApJ...621..215S,2005nfcd.conf..351K, 2002A&A...382...43D, 2000ASPC..218..263V} is another indication that the neutral gas content of low mass galaxies is being altered.  A galaxy moving through intergalactic gas will experience a drag force which, if strong enough, can strip gas from inside the galaxy potential.  The Gunn and Gott condition for a spiral galaxy balances the ram-pressure with the gravitational force per unit area in the plane of the galaxy from stars and gas \\citep{1972ApJ...176....1G}, however, it does not account for the extra gravitational restoring force from the dark matter in the galaxy.  We use a slightly modified form \\citep{2009ApJ...696..385G,2003AJ....125.1926G}, \\begin{equation} 3v^2_{gal} n_{\\mbox{\\tiny IGM}}\\sim \\sigma_*^2 n_{gas} \\end{equation} where $n_{\\mbox{\\tiny IGM}}$ is the density of the intragroup medium, $v_{gal}$ is the velocity of the dwarf galaxy, $n_{gas}$ is the central \\hi density of the dwarf galaxy, and $\\sigma_*$ is the central stellar velocity dispersion in the dwarf.  When the left hand side of the equation becomes larger than the right hand side gas is removed from the galaxy.  It is assumed that the stripping is instantaneous and that the intergalactic gas density encountered by the moving galaxy remains constant.  The use of this criterion as a relatively good approximation is supported by numerical models of gas stripping in clusters \\citep{2000ApJ...538..559M}.  Using the IGM densities measured here we can explore whether ram-pressure stripping in groups is strong enough to remove significant reservoirs of neutral interstellar gas from dwarf galaxies, a key step in producing dSphs.  To model dwarf galaxies we choose the range of total mass, \\hi mass, $\\sigma_v$, and interstellar neutral gas density shown in Table \\ref{tab:dwarfs}.  We compute the combination of galaxy velocity and IGM density necessary to instantaneously strip the neutral gas from a given dwarf galaxy using Equation 3 and show these curves in Figure \\ref{fig:dwarfs} for the three dwarfs differentiating them by \\hi mass.  The density and velocity characteristics of four groups from this paper are plotted to show which will have the necessary conditions to strip \\hi from the dwarfs.  Comparing these stripping curves for generic dwarf galaxies to the IGM densities and average velocities of groups in this paper, we see that the dwarfs with $M_{HI}=10^5\\ \\msol$ ($M_{tot}=10^6\\ \\msol$) and $M_{HI}=10^6\\ \\msol$ ($M_{tot}=10^7\\ \\msol$) are likely to be stripped in all four groups shown. There are many additional factors that could increase the effectiveness of ram-pressure stripping including feedback from star-formation, the cosmic ionizing background, the orbital path of the galaxy, and a fluctuating density distribution.  The densities we derive here suggest that ram-pressure may be a critically important process affecting the gas content of dwarf galaxies.  The question of whether these dwarfs would be stripped in the Local or nearby groups with confirmed morphology-density relations is slightly more complicated.  The Local Group is an order of magnitude less massive ($\\sim 1-2 \\times 10^{12}\\ \\msol$; \\citet{2005AJ....129..178K}) and thus not directly comparable with the systems that we consider here.  We can consider Leo T, a dwarf galaxy in the Local Group which still has a significant amount of \\hi ($\\sim 4 \\times 10^5\\ \\msol$) and is at a Galactocentric distance of 420 kpc.  Leo T is likely to be similar to the progenitors of the newly discovered Milky Way satellites \\citep{2010AdAst2010E..21W} all of which are at radii $< 250$ kpc and have no detected \\hi.  \\citet{2009ApJ...696..385G} estimate that Leo T would have to experience a halo gas density of $0.6-2.3 \\times 10^{-3} \\cmc$ to entirely remove its current \\hi content.  We consider ourselves to be probing intergalactic gas using the radio sources in this paper, however, for systems like the Local Group the dividing line between halo gas in individual galaxies and IGM gas is unclear.   That dwarf galaxies in the Local Group could experience densities in this range is not impossible considering the IGM densities that we report here and the crossing time ($\\sim 2$ Gyr) for the Milky Way system.  In order to produce dSphs the neutral gas must be removed and the distribution of stars must also change.  Tidal stirring caused by repeated shocking at the pericenter of the dwarf's orbit can dissipate angular momentum, transforming the rotationally supported stellar distribution into one that is pressure supported \\citep{2006MNRAS.369.1021M,2011arXiv1104.4278M}.  In these simulations, tidal heating increases the effectiveness of ram-pressure stripping by expanding the dwarf galaxy mass distribution.  Ram pressure stripping is likely not strong enough to remove large quantities of neutral gas from galaxies with total masses of $\\sim 10^9 - 10^{10}\\ \\msol$ which is where the \\himf for groups of galaxies begins to flatten out.  In this regime, it may be more likely that these larger galaxies experience strong tidal stripping through interactions with neighboring galaxies similar to the interaction between the Large and Small Magellanic Clouds and the Milky Way.  In the group environment the space velocities are similar to the internal rotational velocities for large galaxies which makes tidal interactions especially damaging.  \\hi observations of galaxy groups show numerous examples of intergalactic \\hi that appears to be tidally stripped or large galaxies which are \\hi deficient for their morphological type \\citep{1994Natur.372..530Y,2001A&A...377..812V,2009AJ....138..295F}.     \\begin{figure} \\plotone{f7.eps} \\caption{A comparison of the baryon fraction for group S4 with binned data for groups and clusters from \\citet{2009ApJ...703..982G}.  The mass for the S4 group is a dynamical mass derived from the velocity dispersion while the binned data are X-ray derived system masses.  For S4 the error bars reflect the range in dynamical mass (given the range in velocity dispersion) and the range in baryon fraction (given lower limits on the IGM gas mass and stellar mass in galaxies).  Note that the binned data are X-ray derived system masses and our S4 group mass is a dynamical mass.  We have not included a contribution from intragroup stars.  S4 is located at a radius of $700$ kpc from the group center and is probing gas with a density of $50\\rho_{crit}$.  \\textit{Galaxy groups contain significant reservoirs of baryons in their intragroup medium.} }\\label{fig:missbar} \\end{figure}   \\subsection{Strangulation} The stripping of hot gas from the halos of galaxies preventing this gas from providing a fuel source for further star-formation is referred to as strangulation or starvation.  There are a handful of observations of X-ray tails from galaxies in a group environment \\citep{2008ApJ...679.1162J,2007ApJ...664..804M,2006MNRAS.370..453R,2005ApJ...630..280M,2004ApJ...617..262S}.  Most simulations focus on ram-pressure stripping of disk and halo gas from galaxies in the cluster environment; there are few that consider low mass groups.  \\citet{2008ApJ...672L.103K} perform a cosmological chemodynamical simulation of a $M \\sim 3 \\times 10^{11}\\ \\msol$ total mass galaxy in a group with $M \\sim 8 \\times 10^{12}\\ \\msol$ whose $\\rho_{\\tiny IGM}$ ranges from $2-5 \\times 10^{-28} \\gcmc$ ($2-5 \\times 10^{-4} \\cmc$, similar to the densities presented here).  They find that the star-formation rate in this galaxy decreases because most of the hot gas is stripped and cannot cool and form stars.  \\citet{2009MNRAS.399.2221B} indicate that ram-pressure stripping of $> 50\\%$ of the halo gas can occur in groups if $\\rho_{\\mbox{\\tiny IGM}} \\sim 7 \\times 10^{-3} \\cmc$ which is likely only the case in the very core of these systems.  X-ray observations of galaxies in groups find that $80\\%$ of the more massive galaxies ($L_K > L_*$) retain, to some extent, a hot gas halo although they appear to be more faint than field galaxies \\citep{2008ApJ...679.1162J}.  They do not report the extent of these X-ray halos although the handful of galaxies which show X-ray tails or asymmetries have lengths that range from $6-50$ kpc.  \\citet{2008MNRAS.383..593M} present an analytical stripping condition that describes what they see in their hydrodynamic simulations of stripping in clusters.  For a Milky Way sized galaxy ($M_{tot}\\sim 10^{12}\\ \\msol$) moving with a velocity of $300 \\kms$ through gas that is the same or ten times the density of its halo gas the halo will be stripped to a radius of $75$ kpc or $7.5$ kpc, respectively.  In the analytical model of \\citet{2006ApJ...647..910H} the hot gas halo is easily stripped from galaxies in the group environment.  More observations are necessary to establish and compare the extent and brightness of hot halos in field and group galaxies.    \\subsection{Baryon Content}  The baryon deficit in the local universe appears to scale with potential well depth such that the baryon content of massive clusters is nearly as expected while groups and individual galaxies are lacking \\citep{2003ApJ...585L.117B, 2010ApJ...719..119D}.  There are recent indications that the possible reservoir of undetected gas in halos around galaxies is not sufficient to make up the balance \\citep{2010ApJ...714..320A}.  X-ray observations of intergalactic gas in clusters and more massive groups of galaxies are able to trace the radial density distribution of the hot IGM.  However, the current generation of X-ray telescopes are not able to detect emission from intergalactic gas in low mass galaxy groups or emission at large radii in more massive systems, likely because of the temperature of this gas.  In F08 we put an upper limit of $2 \\times 10^6$ K on the temperature of the IGM at the location of bent-double radio source S1 using our density measurement and a non-detection in the X-ray data; that upper limit holds under the analysis in this paper.  With this method we are able to probe the gas density at a single location in a given group but would still like to estimate the total mass in intergalactic gas.  In what follows we make an estimate of this mass and the baryon fraction of the S4 group.  We choose the S4 group for this exercise because its properties are most similar to the groups studied in \\citet{2009ApJ...703..982G} and, in general, it's best to probe the baryon fraction at the largest possible radius.  The location of S4 at $R_{group}=700$ kpc allows us to probe the system at twice the radius possible for the S1 or S3 groups.  By assuming a uniform density within the radius of the bent-double radio source we put a lower limit on the mass in intergalactic gas of $M_{\\mbox{\\tiny IGM}}> 1 \\times 10^{13}\\ \\msol$.  Using the velocity dispersion and its $68\\%$ confidence intervals we estimate a range in dynamical mass for the S4 group of $M_{dyn} = 5_{-1}^{+2}\\times 10^{13}\\ \\msol$.  We calculate a lower limit on the stellar mass of group galaxies ($M_* > 1.3 \\times 10^{12}\\ \\msol$) by using the MPA/JHU value-added DR7 SDSS catalog of stellar masses (determined from photometry) as only the four brightest group members are present in the catalog.  In Figure \\ref{fig:missbar} we compare the baryon fraction estimated above for the S4 group with the COSMOS study of baryon mass fractions in groups and clusters with $z \\leq 1$ \\citep{2009ApJ...703..982G}.  They compute the baryon fraction at $R_{500}$ ($\\gtrsim 500$ kpc for most groups) including contributions from intergalactic gas (from a fit to the X-ray gas mass fraction for 41 well-detected systems), galaxy stellar masses, and intergalactic stars.  Their total mass estimate for a system comes from an $L_X-M_{200}$ relation established by way of a weak lensing analysis.  The $\\Lambda$CDM cosmological baryon fraction by mass is $0.18$ \\citep{2011ApJS..192...18K} assuming a Hubble constant of $73 \\kmsmpc$.  For source S4, located at a radius of 700 kpc this corresponds to a range in baryon fraction of $0.24^{+0.09}_{-0.08}$ which is consistent with what is expected cosmologically although we have not taken into account the errors on the IGM density.  While we are unable to tightly constrain the baryon content of the intragroup medium in these systems, the gas densities that we measure here are valuable as complementary, model and temperature independent, measurements to X-ray and UV absorption line observations of the IGM in galaxy groups.  They indicate intergalactic gas does exist in galaxy groups, even at large radii, and can account for a large fraction of the baryons which are thought to be missing from these systems."
    },
    "1107/1107.0154_arXiv.txt": {
        "abstract": "We perform calculations of isolated disc galaxies to investigate how the properties of the ISM, the nature of molecular clouds, and the global star formation rate depend on the level of stellar feedback. We adopt a simple physical model, which includes a galactic potential, a standard cooling and heating prescription of the ISM, and self gravity of the gas. Stellar feedback is implemented by injecting energy into dense, gravitationally collapsing gas, but is independent of the Schmidt-Kennicutt relation. We obtain fractions of gas, and filling factors for different phases of the ISM in reasonable ageement with observations. Supernovae are found to be vital to reproduce the scale heights of the different components of the ISM, and velocity dispersions. The GMCs formed in the simulations display mass spectra similar to the observations, their normalisation dependent on the level of feedback. We find $\\sim$40 per cent of the clouds exhibit retrograde rotation, induced by cloud--cloud collisions. The star formation rates we obtain are in good agreement with the observed Schmidt-Kennicutt relation, and are not strongly dependent on the star formation efficiency we assume, being largely self regulated by the feedback. We also investigate the effect of spiral structure by comparing calculations with and without the spiral component of the potential. The main difference with a spiral potential is that more massive GMCs are able to accumulate in the spiral arms. Thus we are able to reproduce massive GMCs, and the spurs seen in many grand design galaxies, even with stellar feedback. The presence of the spiral potential does not have an explicit effect on the star formation rate, but can increase the star formation rate indirectly by enabling the formation of long-lived, strongly bound clouds. ",
        "introduction": "Star formation in galaxies is intrinsically linked to the evolution and dynamics of the interstellar medium and in particular, molecular clouds. Thus providing a plausible model of the evolution of galaxies, and star formation within them requires being able to match the properties of the ISM and molecular clouds, as well as reproducing a star formation rate in agreement with the Schmidt-Kennicutt \\citep{Schmidt1959,Kennicutt1998} relation.  In a previous paper \\citep{Dobbs2011}, we examined one particular aspect of molecular clouds. We demonstrated that with a simple model involving a standard cooling prescription, UV heating and feedback from star formation, it was possible to produce clouds with a distribution of virial parameters close to that observed, and that through cloud-cloud collisions, and supernovae feedback, the majority of the clouds were unbound. In the current paper we consider many more aspects of our models, including the global properties of the ISM, star formation rates, and the importance of spiral structure, with an overall aim to match the observed properties of galaxies, as described below.  The ISM in galaxies is characterised by a multiphase turbulent medium, regulated by self gravity, cloud collisions, supernovae and stellar winds, galactic rotation, spiral shocks and magnetic fields. These processes lead to a velocity dispersion of 5--10~km~s$^{-1}$ \\citep{derKruit1982,Dickey1990,Dickey1990a,Combes1997,vanZee1999,Petric2007, Tamburro2009,Wilson2010} both in HI and CO, in fact in all but the very hot gas. The main source of this velocity dispersion however is not fully clear, although the recent consensus is that turbulence is driven on large scales, e.g. by supernovae or spiral shocks \\citep{Ossenkopf2002,Brunt2003,Brunt2009,Padoan2009}. The thermal distribution of the ISM is not well constrained but observations of the solar neighbourhood suggest that similar proportions of HI lie in the cold, warm and intermediate regimes \\citep{Heiles2003}. A potentially difficult characteristic of the ISM to reproduce is the scale height of the gas, particularly the cold phase. The presence of high latitute GMCs such as Orion requires that molecular, and cold HI gas resides at distances above the mid-plane much higher than expected simply from thermal pressure. Previous simulations showed that in the absence of stellar feedback, the scale height of the gas is too low, and the distribution of HI in the $l-z$ plane does not match the observations \\citep{Douglas2010}.  In previous calculations, we showed that giant molecular clouds (GMCs) form in spiral galaxies when smaller clouds are brought together in the spiral arms \\citep{Dobbs2008}. Self gravity aids this process predominantly by increasing the frequency of collisions, and the likelihood that clouds merge, due to the mutual gravitational attraction of the clouds. An immediate question is whether GMCs can still form by this means with stellar feedback, or whether feedback disrupts smaller clouds before more massive GMCs can form \\citep{Ostriker2007}. In addition, the previous calculations without stellar feedback found cloud mass functions in agreement with observations, and further showed that retrograde clouds are naturally reproduced as a results of cloud collisions \\citep{Dobbs2008}. However these properties, at least the cloud mass spectrum, are likely to be dependent on stellar feedback.  Observationally, the star formation rate is correlated with the global surface density according to the Schmidt-Kennicutt relation. There is considerable scatter in the star formation rate ($\\gtrsim$ 1 order of magnitude) for a given surface density, whilst the slope of the relation is found to vary according to different gas tracers, from 1.0 for dense tracers such as CO and HCN \\citep{Wong2002,Gao2004,Bigiel2008,Genzel2010}, to 1.4 for the total gas surface density \\citep{Kennicutt1998,Wong2002} but with a steeper turnover below $\\sim 10$ M$_{\\odot}$ pc$^{-2}$. Nevertheless these observations provide our best guide for linking galactic scale physics to small scale star formation.  A fundamental question in relation to the properties of the ISM and molecular clouds, and the star formation rates in galaxies is the importance of spiral structure. \\citet{Roberts1969} proposed that spiral density waves could trigger star formation in the spiral arms. Thus the presence of a spiral density wave could explain some of the scatter observed in the Schmidt-Kennicutt relation. However whether spiral shocks trigger star formation has been strongly debated in the past. \\citet{Elmegreen1986} examined star formation rates in galaxies with different Hubble types, including flocculent and grand design galaxies. Finding no variation in the star formation rate for the different galaxy types, they concluded that density wave triggering was only a small contribution to the star formation rate. The main evidence of density wave triggered star formation rates is a non-linear dependence of the star formation rate in the spiral arms \\citep{Cepa1990,Seigar2002}.  More recent observations however again suggest that there is little difference in the star formation efficiency in the spiral arms \\citep{Foyle2010}. Thus the density wave may simply organise the gas and star formation within a galaxy \\citep{Tan2010}. In \\citet{Dobbs2009} we found that there was little dependence of the amount of bound gas on the strength of the spiral potential, mainly because although the density increases in the spiral shock, the velocity dispersion also increases. In \\citet{Dobbs2009}, the clouds are also supported by magnetic fields.  Finally, whilst star formation occurs predominantly in the spiral arms, many galaxies (typically those with well defined dust lanes) exhibit star formation associated with interarm spurs \\citep{LaVigne2006}. Calculations without feedback \\citep{Kim2003,Wada2004, Shetty2006,Dobbs2007} show that spurs form by the shearing out of spiral arm GMCs. However it is unclear whether these structures would still form if the clouds are disrupted by stellar feedback.  Simulations of isolated galaxies designed to model the ISM with supernovae feedback were performed by \\citet{Rosen1995}, and later by \\citet{Wada2000}. Both demonstrated that a 3-phase medium is produced, and the higher resolution of the latter studies indicated the complex structure of the ISM. However both were 2D, so unable to calculate properties of molecular clouds. The difficulty of performing high resolutions of a 3D disc has lead to many simulations of kpc or so size periodic boxes of the ISM \\citep{Korpi1999,deAvillez2000,deAvillez2004,deAvillez2005,Slyz2005}. These show that supernovae feedback can reproduce the observed levels of turbulence \\citep{Dib2006,Joung2006} and pressure \\citep{Joung2009,Koyama2009} in the ISM. One other result from these calculations is that supernovae feedback appears to inhibit, rather than enhance, star formation \\citep{Slyz2005,Joung2006}. These calculations are able to resolve turbulence and cooling instabilities on parsec scales. However they cannot examine processes such as molecular cloud formation by cloud collisions, or self gravity, which operate on larger scales.  Numerous calculations have modelled the ISM on galactic scales with the aim of reproducing the Schmidt-Kennicutt relation \\citep{Li2006,Tasker2008b,Robertson2008,Pelupessy2009,Gnedin2009}. \\citet{Gnedin2009} also include a prescription of molecular gas formation which they use to investigate different metallicity environments. One should note though that many calculations implicitly assume (or explicitly in the case of cosmological simulations) some form of the Schmidt-Kennicutt relation \\citep{Schmidt1959,Kennicutt1989} in adopting a star formation rate per unit time. Other calculations explicitly insert the observed supernovae rate in our Galaxy \\citep{deAvillez2004,Shetty2008}.  A couple of calculations have also included a spiral potential and stellar feedback in simulations of isolated discs \\citep{Shetty2008,Wada2008}. These tend to show that feedback is fairly disruptive, in the case of \\citet{Shetty2008} even destroying the spiral structure, thus they have difficulty reproducing the spurs observed in galaxies.  Our aim in this paper is to provide a coherent, and global, description of the dynamics of the ISM, with relatively simple physical processes and a minimum of assumptions. The paper is organised as follows. We first describe how the calculations are set up in Section 2. We then provide a generic picture of the typical evolution of our calculations in Section 3. In Section~4 we present simulations where we investigate the structure of the disc and properties of the ISM for calculations with different star formation efficiencies, all with a spiral potential. In Section~5, we compare results from calculations with and without a spiral potential. In Section~6, we discuss results from higher surface density calculations, and show the dependence of the global star formation rate on surface density. We calculate the properties of GMCs from a selection of the calculations in Section~7. Finally in Section~8 we present our conclusions. ",
        "conclusions": "In order to model the ISM in spiral galaxies we have carried out numerical three-dimensional hydrodynamic simulations of gas flowing in a fixed, full disc, global potential. The model involves a standard cooling law, and heating from background UV radiation together with localised input from star-formation episodes. Star formation is assumed to occur instantaneously when a parcel of the ISM becomes sufficiently compact and self-gravitating that it undergoes dynamical collapse. When this occurs a fraction, $\\epsilon$, of this gas is assumed to form stars and to provide instantaneous feedback to the ISM. Using these simple ideas we are able to reproduce the following features of the ISM:  \\subsection{Gas properties} \\begin{enumerate} \\item{We obtain roughly equal fractions of gas which are in the cold ($T < 150K$), intermediate, thermally unstable ($150K < T < 5000K$) and warm ($T > 5000K$) phases. The relative fractions depend mainly on the assumed efficiency $\\epsilon$ of star- formation feedback (Figure 4).} \\item{The scaleheights of the cold, unstable and warm phases are in reasonable agreement with observations. Here values of $\\epsilon$ in the range 0.1 -- 0.2 give the best fit to the observations of the Galaxy (Figure 7), as well as for a number of external early-type spirals (Figure 8).} \\item{The velocity dispersion both in the plane, $\\sigma_r$, and perpendicular to it, $\\sigma_z$ are similar to the observed velocity dispersions. Values obtained range from an average of around 6 km s$^{-1}$ for $\\epsilon$ = 0.05 to around 12 km s$^{-1}$ for $\\epsilon$ = 0.2 (Figure 5). For the smaller values of $\\epsilon$ we find that $\\sigma_r \\approx \\sigma_z$ , and note that values in the interarm regions are marginally smaller ($\\approx$ 5 km s$^{-1}$) than in the arms ($\\approx$ 7 km s$^{-1}$). For the higher values there is little arm/interarm dependence, but there is then a slight trend for $\\sigma_z$ to be on average larger than $\\sigma_r.$} \\end{enumerate} \\subsection{Molecular cloud properties} \\begin{enumerate} \\item{The mass spectra of clouds are of the form $dN/dM \\propto M^{\u22121.9\\pm0.1}$ for values of $\\epsilon$ in the range $0.05 \\leq \\epsilon \\leq 0.2$, with the masses being shifted to lower values for the higher values of $\\epsilon$ (Figure 20).} \\item{The range of values of the $\\alpha$ parameter which describes the degree to which the clouds are bound (see also \\citealt{Dobbs2011}). For a surface density of $\\Sigma = 8M_{\\odot}$ pc$^{-2}$, and for values of $\\epsilon$ in the range $\\epsilon=$ 0.05 -- 0.2 it is found that most clouds are unbound, or marginally bound, and have $\\alpha$ in the range $\\alpha \\approx$ 1.0 -- 10. For smaller values of $\\epsilon$ the models produce clouds which are disproportionately massive and strongly gravitationally bound. Higher surface density discs also tend to show more bound clouds for equivalent values of $\\epsilon$.} \\item{ If molecular clouds formed simply by gravitational collapse from the ISM then conservation of angular momentum acquired from the galactic shear would ensure that they all rotate in the prograde direction. In fact, both in the Galaxy and in M33 about 40 per cent of clouds rotate in the retrograde direction. As we have remarked before \\citep{Dobbs2008}, if clouds build predominantly by collisional build-up of dense structures within an already inhomogeneous ISM, then the tendency of clouds to display both prograde and retrograde rotations in almost equal measure can be explained. We confirm this result in the current simulations (Figures 21 and 22). In Figure 22 we note that for the models with no spiral structure, in which the clouds build preferentially by self-gravity rather than collisions, the fraction of retrograde clouds is significantly reduced.} \\end{enumerate}  \\subsection{Star formation} \\begin{enumerate} \\item{After initial transient behaviour, we find for $\\epsilon \\approx$ 0.05--0.2 that global star formation rates settle down to equilibrium values (Figures 11, 13 and 17) which depend on the average galactic gas surface density in agreement with the observed findings of Kennicutt (2008) (Figure 18). This result is independent of the presence or absence of an imposed spiral potential (cf. \\citealt{Dobbs2009}). For efficiencies of $\\epsilon \\lesssim 0.5$ we find that the spiral potential gives rise to long-lived, massive, bound clouds in which star formation continues to accelerate over a long period; for these calculations we are unable to determine an equlibrium star formation rate.} \\item{We find that for an imposed spiral potential the distance from the galactic plane where star formation occurs depends significantly on the assumed value of $\\epsilon$ (Figure 9). For the lower value of $\\epsilon=0.05$ the star formation events occur with a scaleheight of 40 pc and do not stray more than around 150 pc above and below the galactic plane. In contrast, for $\\epsilon=0.2$ the scaleheight of star formation events is around 80 pc, and events can occur up to around 500 pc from the plane.} \\end{enumerate}  In addition to these properties, we also obtain notable disagreements with the observations when there is little or no feedback, as exemplified by our calculation where $\\epsilon=$ 0.01. In this case, the mass spectrum is bimodal, there are very few retrograde clouds, whilst the scale height of the ISM is too low in the absence of feedback."
    },
    "1107/1107.0224_arXiv.txt": {
        "abstract": "{The nearly edge-on galaxy NGC~4945 is one of the closest galaxies where an AGN and starburst coexist, and is one of the brightest sources at $100~\\rm keV$. Near and mid-infrared spectroscopy have shown very strong obscuration of its central region, rivaled only in strength by some of the most deeply obscured ULIRGs. } { Determine the spatial distribution of ISM emission features in the central 426$\\times$426 $\\rm pc^2$ of NGC~4945. } {We map the central region of NGC~4945 in three of the four Spitzer-IRS modules (SH, SL and LL). In particular, we produce maps of the flux distribution of the starburst tracers \\neii, \\neiii, \\siii\\ and \\siv\\ at 12.81, 15.56, 18.71 and 10.51 \\mum, respectively, and a map of the AGN narrow-line region tracer \\nev\\ at 14.32 \\mum. In addition, we mapped the S(1), S(2) and S(3) pure rotational lines of H$_2$, which trace the distribution of warm molecular hydrogen. Finally, we obtained an extinction map ($A_{\\rm V}$) based on the apparent strength of the 9.7~\\mum\\ silicate absorption feature.} { At a spatial resolution of $\\sim$5$''$ our extinction map traces the contours of the starburst ring. The highest extinction is, however, found at the nucleus, where we measure $A_{\\rm V}(9.85~\\mu \\rm m)$$\\approx$60. Within the uncertainty of the astrometry all emission lines are found to peak on the nucleus, except for the warm molecular hydrogen emission which shows a maximum 60-100 pc NW of the nucleus. We favour a scenario in which the lower H$_2$ 0-0 S(1) and S(2) rotational lines originate mainly from an unobscured extra-nuclear component associated with the super-wind cone observed in the HST NICMOS map of the H$_2$ 1-0 S(1) vibrational line. For the \\nev\\ emission we infer an attenuation of a factor 12-160 ($A_{\\rm V}$=55-112) based on a comparison of the ratio of our \\nev\\ flux and the absorption-corrected 14--195 keV Swift-BAT flux to the average \\nev/BAT ratio for Seyfert 1 nuclei. The high attenuation indicates that \\nev\\ and \\oiv\\ cannot be used as extinction-free tracers of AGN power in galaxies with deeply buried nuclei.} {} ",
        "introduction": "\\label{sec:intro}  At a distance of $\\sim$3.82 Mpc ($\\sim$18.5 pc/arcsec, \\citealt{karachentsev07}) the active galaxy NGC~4945 is one of the closest galaxies that host both, an AGN and starburst. Earlier X-ray observations showed evidence for a hidden AGN \\citep{iwasawa93, guainazzi00}. These observations revealed a Compton-thick spectrum with an absorbing column density of $N_{\\rm H}=2.4\\times 10^{24}~\\2cm$ \\citep{guainazzi00}. The nucleus of NGC~4945 is one of the brightest extragalactic sources at $100~\\rm keV$ \\citep{done96}, and the brightest Seyfert 2 AGN at $>20~\\rm keV$ \\citep{itoh08}.  Rather than a point source marking the presence of the AGN in optical and near-infrared images, the nearly edge-on ($i\\sim 80^o$) line of sight to the central region reveals evidence for strong and patchy extinction, which is especially apparent in HST-NICMOS $H-K$ maps \\citep{marconi00}. A dust lane aligned along the major axis of the galactic disk obscures parts of the central region just southeast of the $K$-band peak. The $K$-band peak itself lies $\\sim1''$ west (but, within the uncertainty) of the position of the H$_2$O mega maser \\citep{greenhill97}, which we adopt as the location of the AGN. The mega maser emission allows to estimate the mass of the central black hole, $M_{\\rm BH}\\approx1.4\\times10^6~\\rm M_{\\sun}$, which together with its inferred 0.1--200 keV bolometric luminosity  of $\\sim$$2\\times10^{43}~\\ergs$ ($\\sim5\\times10^9~L_{\\sun}$, \\citet{guainazzi00}) indicates that the central engine radiates at $\\sim$10\\% of its Edington luminosity \\citep{greenhill97, madejski00}.  Estimates from IRAS observations indicate that about 75\\% of the total infrared luminosity of the galaxy (L$_{\\rm IR}=2.4\\times10^{10} ~L_{\\sun}$) is generated within an elongated region of $12''\\times9''$ (about $222\\times167~\\rm pc^2$) centered on the nucleus \\citep{brock88}. The structure of this region, as shown in high resolution HST-NICMOS observations of the Pa$\\alpha$ line, is consistent with a nearly edge-on starburst disk with a $4.5''$ ($\\sim83~\\rm pc$) radius \\citep{marconi00}, embedded in a $r=8''$ ($\\sim$150 pc) inclined (62 degrees) molecular disk \\citep{chou07}.  Although the star formation and supernova rates in the nuclear region were originally estimated to be moderate ($\\sim0.4$ \\Msun\\ $\\rm yr^{-1}$ and $\\sim0.05~\\rm yr^{-1}$, respectively; \\citealt{moorwood94}), more recent estimates based on high resolution (angular scale of $0.3~\\rm pc$) radio observations and estimates of supernova remnant source counts, sizes and expansion rates, lead to a type II supernova rate of $>0.1(v/10^4)~\\rm yr^{-1}$, and star formation rate (SFR) limits of $2.4(v/10^4)<{\\rm SFR}(M\\geq5~M_{\\sun}) < 370~M_{\\sun}~\\rm yr^{-1}$, where $v$ is the shell radial expansion velocity in \\kms\\ \\citep{lenc09}. These supernova and star formation rates are, within a factor of two, similar to those estimated in NGC~253 and M82 \\citep{pedlar03, lenc06}, which are also nearly edge-on starburst galaxies, have similar distances close to $4~\\rm Mpc$ and all show similar infrared luminosities \\citep{rice88}. The impact of the central starburst on the circumnuclear region is large, as revealed by the presence of a conical cavity evacuated by a supernova driven wind \\citep{moorwood96a, marconi00}. The HST-NICMOS images of the H$_2$ 1-0 S(1) line at $2.12$ \\mum\\ show that the edges of the cavity extend out to $5''$ ($\\sim$93$\\rm pc$) north from the NGC~4945 nucleus \\citep{marconi00}.  Near and mid-infrared spectroscopy indicate that the ISM line of sight to the central region of NGC~4945 is very different than that of other nearby starburst galaxies, such as M~82 and NGC~253, which present similar inclination angles. The strong absorption features of both volatile (CO and CO$_2$) and refractory (H$_2$O) ices observed in the 2.4--5 \\mum\\ ISO-PHT-S spectrum indicate the presence of shielded cold molecular clouds obscuring the NGC~4945 nucleus \\citep{spoon00, spoon03}. The detection of a strong absorption feature of XCN ice at 4.62~\\mum\\ by \\citet{spoon00, spoon03} indicates that these molecular clouds have been processed in an energetic environment \\citep{lacy84}, and processed ice is suggested to be a common characteristic of dense molecular material in star forming galactic nuclei \\citep{spoon03}. A highly distorted PAH emission spectrum produced by a very deep 9.7 \\mum\\ silicate absorption feature is more evidence for the unusually strongly obscured nuclear region \\citep{spoon00, brandl06}. This makes NGC~4945 a unique nearby laboratory to study an environment that can be found only in distant ULIRGs.   Previous mid-infrared spectroscopic observations could not confirm the presence of an AGN in NGC~4945. Just an upper limit of the 14.32 \\mum\\ \\nev\\ emission, considered a tracer of AGN narrow-line regions \\citep[e.g.,][]{moorwood96b, genzel98}, was obtained from ISO-SWS observations \\citep{spoon00}. The VLT-ISAAC observation of the 3.93 \\mum\\ \\SiIX\\ line, commonly observed in the soft X-ray photoionized gas of many Seyfert galaxies \\citep{oliva94, lutz02}, resulted in a non-detection \\citep{spoon03}. Only recent observations with the more sensitive IRS-SH spectrograph on Spitzer allowed the detection (although at a very faint level) of the 14.32 \\mum\\ \\nev\\ line towards the NGC~4945 nucleus \\citep{bernards09}. The \\nev14.3\\mum/\\neii12.8\\mum\\ flux ratio found in NGC~4945 is $\\sim0.007$, which indicates an AGN contribution of less than a few percent using the diagram by \\citet{farrah07} (their Fig. 16). This detection of \\nev\\ relative to \\neii\\ is about 10 times weaker than what was observed in other Seyfert galaxies like Mrk~266 and NGC~1365 \\citep{bernards09}.  In this paper we study the mid-IR properties of the nuclear region of NGC~4945. Spitzer-IRS spectral mapping observations of a $23''\\times23''$ (about 426$\\times$426$~\\rm pc^2$) region are presented. The mapping capabilities of Spitzer-IRS allow the study of a number of properties of the nuclear region in NGC~4945. For instance, the extent of the AGN coronal line region as traced by the \\nev\\ line, the extent of the crystalline silicate absorbing region (and a search for the sources responsible for the presence of crystalline silicates), disentangling PAH emission and silicate absorption along the line of sight, and the excitation and age of the circumnuclear-nuclear starburst based on line ratios of forbidden lines. Although we actually used the Spitzer-IRS modules SH, SL and LL, in this work we present the most interesting results from the SH and SL spectral maps only, since we focus on the analysis of the spatial distribution of ISM emission and absorption features. We present for the first time the maps of the starburst tracers \\neii\\ 12.81 \\mum, \\neiii\\ 15.56 \\mum, \\siii\\ 18.71 \\mum, and \\siv\\ 10.51 \\mum, and a map of the AGN narrow-line region tracer \\nev\\ at 14.32 \\mum, and analyze the ratios between some of these tracers as different diagnostics. The organization of this article is as follows. In Sect.~\\ref{sec:obs} we describe the observations and the data reduction. The maps obtained are presented in Sect.~\\ref{sec:results}. The analysis of the data is presented in Sect.~\\ref{sec:discuss}. The conclusions and final remarks are presented in Sect.~\\ref{sec:final-remarks}. ",
        "conclusions": "\\label{sec:discuss}    \\subsection{Tracing the circumnuclear starburst ring}\\label{sec:starburst-ring}  The \\neii\\ and \\siii\\ fine-structure lines are our cleanest tracers of star formation (\\hii\\ regions). If an AGN is present, however, the \\neiii\\ and \\siv\\ lines will have contributions from both the starburst and the AGN. In NGC~4945 the AGN contamination to \\neiii\\ is likely low, given the faintness of \\nev. In fact, at the peak flux of the \\nev\\ and \\neiii\\ maps (Fig.~\\ref{fig:SH-maps1}) the former corresponds to only about 5\\% of the \\neiii\\ flux.  From Table~\\ref{tab-c4:fluxes-fov} we derive the respective luminosities of $L_{\\rm [Ne V]}=7.42\\times10^{29}~\\Wsr$ and $L_{\\rm [Ne III]}=1.76\\times10^{31}~\\Wsr$, using a luminosity distance $D_L$=3.82 Mpc. However, from the relation between $L_{\\rm [Ne V]}$ and $L_{\\rm [Ne III]}$ by \\citet[][]{gorjian07} (which holds for a group of Seyfert 1 and 2, AGNs, ULIRGs and radio galaxies), the predicted \\neiii\\ luminosity should be $\\sim3.28\\times10^{30}~\\Wsr$, that is a factor $\\sim$5.4 lower than observed. This means that NGC~4945 has a slightly lower AGN contribution to the \\neiii\\ line than the outlier NGC~3079 found in the sample by \\citet[][]{gorjian07}, for which the \\neiii\\ luminosity is a factor $\\sim$5 brighter than expected based on its \\nev\\ luminosity.        \\begin{figure*}[!htp]      \\hfill\\includegraphics[width=14.2cm]{fig13}\\hspace*{\\fill}   \\caption{\\footnotesize{\\textit{Left panels} - Integrated flux ratios between the fine-structure lines (from top to bottom) \\nev~14.32/\\neii~12.81, \\neiii~15.56/\\neii~12.81, \\siii~18.71/\\neii~12.81, and \\siv~10.51/\\siii~18.71. \\textit{Right panels} - Same as above, but considering the fluxes corrected for extinction. The hatched areas correspond to pixels with a $<$3$\\sigma$ detection level in the faintest emission lines \\nev~14.32~\\mum\\ and \\siv~10.51~\\mum.}} \\label{fig:line-ratios} \\end{figure*}     The \\neiii~15.56/\\neii~12.81 \\mum\\ line ratio (Fig.~\\ref{fig:line-ratios}) ranges between $\\sim$0.13 and $\\sim$0.27 above and below the major axis of the region mapped. The difference with the lowest ratios (0.06--0.13), which are found along the northeast-southwest axis, is more pronounced than in the \\nev~14.32/\\neii~12.81 \\mum\\ line ratios (\\textit{top left panel} in Fig.~\\ref{fig:line-ratios}).  The lowest \\neiii~15.55~\\mum/\\neii~12.81~\\mum\\ line ratios coincide with the 100 pc-scale circumnuclear starburst ring of NGC~4945 (\\textit{bottom panels} in Fig.~\\ref{fig:PAalpha}), detected in Pa$\\alpha$ emission with HST NICMOS \\citep{marconi00}. The increasing gradient above the starburst ring seems to mimic the conical-shaped cavity traced by the H$_2$ 1-0 S(1) line \\citep{marconi00}. However, the ratio also increases toward the south-east direction from the starburst ring, which rather implies that the ring has different properties than its surroundings.  In fact, a similar increase of \\neiii/\\neii\\ ratios away from the nucleus is observed in the nuclear maps by \\citet[][their Fig.5]{pereira10}, with exception of NGC~7130, the only galaxy classified as LINER/Seyfert 2 in their sample. According to \\citep{snijders07} the higher the metallicity, the higher the age of the stellar population and the higher the gas density, the lower should be the \\neiii~15.55~\\mum/\\neii~12.81~\\mum\\ ratio.  The \\siii~18.71/\\neii~12.81 line ratio is considered a good tracer of densities in the range $10^4$ and $10^6~\\3cm$ because their similar excitation potentials (21.6 eV and 23.3 eV, for \\neii\\ and \\siii, respectively) and significantly different critical densities ($6.1\\times10^5~\\3cm$ for \\neii\\ and $1.0\\times10^4~\\3cm$ for \\siii) make this ratio less sensitive to the hardness of the radiation field than to the density of the ISM. According to the models by \\citet{snijders07} a lower \\siii/\\neii\\ ratio would indicate larger densities. In Fig.~\\ref{fig:line-ratios} the \\siii~18.71/\\neii~12.81 line ratios are also lower in the center than around it. This ratio ranges between 0.002 and 0.013 in the center, which indicate densities larger than $10^6~\\3cm$ in a $>$5 Myr old starburst system with solar metallicity and relatively high ($\\rm q=8\\times10^8$) ionization parameter \\citep[][their Fig.5]{snijders07}. Based on our observed \\neiii/\\neii\\ ratios, the \\siii/\\neii\\ ratios predicted with the model by \\citet[][their equation 2]{pereira10} are about ten times larger than the observed ratios. These can be a consequence of the about 10 times higher extinction found in NGC~4945 than in the sample of galaxies used by \\citet{pereira10}.  The \\neiii/\\neii\\ line ratios obtained with the extinction correction (\\textit{middle right panel} in Fig.~\\ref{fig:line-ratios}) are just $\\sim$9\\% larger than without correction. This relatively small change after the extinction correction is because even in a high-extinction situation the differential extinction between \\neiii\\ and \\neii\\ is small, given that both lines are closely spaced in wavelength and not in one of the silicate absorption features. On the other hand, because of their larger differential (wavelength) extinction, the \\siii/\\neii\\ ratios corrected for extinction do change significantly from a factor $\\sim$50 in the center (where the extinction is larger) to a factor 3 away from the center (where the extinction is lower).      \\begin{figure}[tp]        \\hfill\\includegraphics[width=6.5cm]{fig14}\\hspace*{\\fill}  \\caption{\\footnotesize{HST NICMOS image of Pa$\\alpha$ (in units of $10^{-21}~\\Wcm$) of the nucleus of NGC~4945 \\citep{marconi00}. The not corrected for extinction contours (labelled in units of $10^{-12}~\\Wcmsr$) are the fine-structure lines \\neii~12.81~\\mum\\ (\\textit{top panel}) and \\nev~14.32~\\mum\\ (\\textit{middle panel}). The \\nev\\ flux peaks at about one pixel to the northwest of the H$_2$O mega maser (Fig.~\\ref{fig:SH-maps1}), The \\textit{bottom panel} shows the \\neiii~15.55/\\neii~12.81~\\mum\\ line ratio not corrected for extinction. The ratio is lower along the starburst ring traced by the Pa$\\alpha$ line, while it increases away from the center.}} \\label{fig:PAalpha} \\end{figure}    Note that only the ratios from our co-added fluxes (Table~\\ref{tab-c4:fluxes-fov}) can be compared to other galactic nuclei, as the 10$\\times$10 aperture is comparable to the size probed in any of the more distant galaxy nuclei. The \\neiii/\\neii\\ ratio at the position of the H$_2$O mega maser is about 10\\% lower than the ratio obtained from the fluxes of the 10$\\times$10 co-added spectrum (Table~\\ref{tab-c4:fluxes-fov}).  Since in starburst environments the \\neii\\ and \\neiii\\ emission lines are expected to be driven mainly by photo-ionization \\citep[e.g.,][]{ho07}, the lower ratios found along the northeast-southwest axis are likely due to a \\neii\\ emission enhanced by the starburst ring. Although, these low ratios are also consistent with a ratio \\neiii/\\neii$\\leq$0.1 found in shocks \\citep{binette85} where the low ionization line \\neii\\ can also be enhanced \\citep{voit92}.  On the other hand, the highest \\neiii/\\neii\\ ratios found above and below the starburst ring are larger than those typically found in shocks. These higher ratios could actually be tracing an additional contribution to the [NeIII] emission by a conically shaped narrow line region (NLR). Previously, no evidence was found for the existence of such a NLR, given the absence of 5007\\AA\\ \\oiii\\ emission in the central 800pc$\\times$800pc \\citep{moorwood96a}. Instead the conical cavity traced by optical and near-infrared line and continuum tracers \\citep{moorwood96a, marconi00} was associated with a starburst super wind. However, given the high extinction implied by our observations, optical \\oiii\\ emission can be easily attenuated; mid-infrared \\neiii\\ emission less so. A hypothetical conical NLR would be difficult to identify from a \\neiii\\ map alone due to blending with the nuclear starburst component. But it may manifest itself more clearly in a \\neiii/\\neii\\ map, in which the contribution of the starburst component diminishes quickly beyond the nucleus proper.   The \\neiii/\\neii$\\sim$0.06--0.13 ratios observed along the starburst ring are in close agreement with the ratios observed in some galaxies of the ISO starburst sample and are consistent with burst timescale models predicting a relatively short-lived starburst of $5-8~\\rm Myr$ \\citep[their Fig.6]{thornley00}. In the whole region mapped, we  observe ratios \\neiii/\\neii$<$0.3, which are a factor $\\sim$3 lower than the ratios observed in the sample of quasars reported by \\citet{veilleux09}. However, the ratio \\neiii/\\neii$\\sim$0.07 observed at the nucleus of NGC~4945 is comparable to the ratios observed in the nucleus of 5 (out of 16) galaxies in the sample of LIRGs studied in \\citet[their Fig.14]{pereira10}.  On the other hand, the differential extinction between \\siii\\ and \\siv\\ is larger than between \\neii\\ and \\neiii. Because \\siv\\ is sitting close to the deepest point of the silicate absorption feature at 9.7~\\mum\\ (Fig.\\ref{fig:average-FOV-spectra}), the \\siv~10.51/\\siii~18.71 \\mum\\ line ratios corrected for extinction (\\textit{bottom right} in Fig.~\\ref{fig:line-ratios}) are a factor $\\sim$2 larger than those obtained without correction (\\textit{bottom left} in Fig.~\\ref{fig:line-ratios}). In both cases the highest ratio is found about $2.3''$ (one pixel) northwest of the H$_2$O mega maser. The uncorrected for extinction map-averaged ratios (from the fluxes in Table~\\ref{tab-c4:fluxes-fov}) of \\neiii/\\neii$\\sim$0.11 and \\siv/\\siii$\\sim$0.02 are similar (within 10\\%) to the ratios obtained from the SH staring observations by \\citet{bernards09}. They position NGC~4945 among the sources with the lowest hardness of the radiation field in their sample of starburst galaxies, which, according to \\citeauthor{bernards09}, may be an indication of an old (or small) population of massive stars in NGC~4945. However, this can be concluded only for the outermost surface that we can probe with \\siv\\ and \\siii\\ lines, and we cannot rely much on the extinction corrected data for this analysis, since the \\siv/\\siii\\ ratio depends more on the extinction law used than the \\neiii/\\neii\\ ratio.         \\subsection{Molecular hydrogen gas}\\label{sec:H2-gas}   \\begin{figure*}[!tp]   \\hfill\\includegraphics[width=14.2cm]{fig15}\\hspace*{\\fill}  \\caption{\\footnotesize{IRS/SH contours of the rotational H$_2$ 0-0 S(1) (\\textit{left}) and H$_2$ 0-0 S(2) (\\textit{right}) lines overlaid on the HST NICMOS map of the vibrational H$_2$ 1-0 S(1) molecular hydrogen line \\citep{marconi00}.}} \\label{fig:HST-H2-IRS-H2} \\end{figure*}  As can be seen in Figures \\ref{fig:SH-maps1} \\& \\ref{fig:SH-maps2}, the emission centroids of the H$_2$ 0-0 S(1), S(2) and S(3) lines are clearly offset from the centroids of the low ionization fine-structure lines, which all peak close to the H$_2$O mega maser. Such an offset is hard to understand if the H$_2$ 0-0 S(1), S(2) and S(3) lines and the HII region gas would be exposed to similar amounts of extinction. The differential extinction between H$_2$ S(2) and \\neii, at 12.28 and 12.81 \\mum, for instance, is very small. The same is true for the wavelengths of H$_2$ S(1) and \\siii\\ at 17.0 and 18.71 \\mum, respectively. The lack of H$_2$ 0-0 S(1), S(2) and S(3) emission at the position of the mega maser can be explained if the molecular hydrogen emission in the nucleus would be more strongly attenuated than the HII region gas. That is, the inner H$_2$ emission would suffer additional/differential extinction that is not accounted for by the silicate feature. On the other hand, analysis of the higher pure rotational lines observed in the SL module (Spoon \\etal\\ \\textit{in preparation}) and with VLT-ISAAC \\citep[][their Figure 7]{spoon03} indicates that the S(5), S(7) and S(9) lines originate progressively closer to the nucleus. This seems hard to reconcile with the above scenario, given that A($\\lambda$)/A$_V$ in the S(5)-S(9) wavelength range is likely higher than in the S(1)-S(2) wavelength range.  We instead favour a scenario in which the lower pure rotational lines S(1) and S(2) mainly originate from an unobscured extra-nuclear component associated with the super-wind cone as seen in the HST NICMOS map of the vibrational H$_2$ 1-0 S(1) emission \\citep{marconi00}. Figure~\\ref{fig:HST-H2-IRS-H2} shows the IRS/SH maps of the pure rotational H$_2$ S(1) and S(2) lines overlaid on the H$_2$ 1-0 S(1) map. Interestingly, the two pure rotational lines peak further from the apex of the cone than the 1-0 S(1) line does. The 0-0 S(1) line (upper level energy 1015 K) more so than the 0-0 S(2) line ($T_{up}$=1682 K). In this scenario the higher pure rotational lines S(5), S(7) and S(9) would originate mainly from the starburst ring or within. There one likely finds the high critical densities ($n_{\\rm crit}>10^5~\\3cm$) as well as high gas temperatures ($T_{\\rm K}>10^3$ K) needed to thermalize these lines \\citep[as seen in active galaxies like NGC~1068, Mrk~231, and Arp~220; e.g. ][]{pb07, pb09, aalto07a, vdwerf10}, but they will also be affected by strong extinction. The latter may cause the nuclear component of the 2.12 $\\mu$m H$_2$ 1-0 S(1) line ($T_{up}$=6950 K) to suffer far more extinction than the longer wavelength H$_2$ 0-0 S(5)-S(9) lines and therefore to be associated solely with the conical cavity. In all of this the H$_2$ 0-0 S(3) line, originating at an upper level energy of 2500 K, may have both strong nuclear and cavity components. Due to the extreme extinction the nuclear component would be suppressed, leaving only the cavity component for us to see.    \\subsection{Excitation of H$_2$}\\label{sec:H2-Tex}  For LTE conditions, and an ortho-to-para ratio of 3 \\citep[their Fig.13]{neufeld06}, we can estimate the excitation temperature of the molecular hydrogen throughout the region mapped with the IRS/SH module, from the ratio between the flux of the H$_2$ 0-0 S(2) 12.3 \\mum\\ and H$_2$ 0-0 S(1) 17.0 \\mum\\ lines as  \\begin{equation}\\label{eq:H2-Tex} T_{ex}=-\\frac{T_{up}^{\\rm S(2)} - T_{up}^{\\rm S(1)}}{ log(F_{\\rm S(2)}\\nu_{\\rm S(1)}A_{\\rm S(1)}g_{\\rm S(1)}) - log(F_{\\rm S(1)}\\nu_{\\rm S(2)}A_{\\rm S(2)}g_{\\rm S(2)}) }~~~\\rm K, \\end{equation}   \\begin{figure*}[!tp]   \\hfill\\includegraphics[width=14.2cm]{fig16}\\hspace*{\\fill}  \\caption{\\footnotesize{\\textit{Left panels} - IRS/SH maps of the H$_2$ 0-0 S(2) to H$_2$ 0-0 S(1) total flux ratio (\\textit{top left panel}), and the map of the estimated $T_{ex}$ ($\\rm K$) of the molecular hydrogen. The peak excitation temperature of $\\sim$528$~\\rm K$ is reached about $\\gtrsim$2.3$''$ (one pixel) to the north of the H$_2$O mega maser. \\textit{Right panels} - Same as above, but using the extinction corrected fluxes of the H$_2$ S(2) and S(1) lines. The extinction correction leads to an $\\sim$100$~\\rm K$ lower peak temperature.}} \\label{fig:H2-Tex} \\end{figure*}  \\noindent where $F_{\\rm S(1)}$ and $F_{\\rm S(2)}$ are the integrated flux densities of the H$_2$ 0-0 S(1) and S(2) lines, $\\nu$ is the corresponding rest frequency ($\\sim1.75\\times10^{13}~\\rm Hz$ for S(1) and $\\sim2.44\\times10^{13}~\\rm Hz$ for S(2)), and $A$ and $g$ are the respective Einstein $A$-coefficients and statistical weights of each transition. The upper level energy $T_{up}$ of the transitions is in units of $\\rm K$.    Figure~\\ref{fig:H2-Tex} shows the IRS/SH maps (not corrected for extinction) of the H$_2$ 0-0 S(2) to S(1) total flux ratio (\\textit{top left panel}), and the corresponding map of the excitation temperature $T_{ex}$ ($\\rm K$) (\\textit{bottom left panel}) of the molecular hydrogen, estimated from eq.(\\ref{eq:H2-Tex}). The peak excitation temperature of $\\sim528\\pm31~\\rm K$ is reached $\\gtrsim$2.3$''$ (one pixel) to the north of the H$_2$O mega maser. The \\textit{right panels} of Fig.~\\ref{fig:H2-Tex} show the H$_2$ 0-0 S(2)/S(1) ratio and the estimated excitation temperature from the respective H$_2$ fluxes corrected for extinction with the silicate-based $A_{\\rm V}(9.85~\\mu \\rm m)$. The correction for extinction leads to a peak temperature of $421~\\rm K$, which is about $100~\\rm K$ lower than the temperature derived from the non-corrected fluxes. However, the distribution of the temperature, and the position of its peak value, do not change.  Since galaxies are complex systems, the H$_2$ gas is not expected to be at a single temperature. Besides, a long line of sight can probe different excitation environments (e.g., PDRs, XDRs, shocks) like in the sample of galaxies studied by \\citet{roussel07}. Previous estimates of the excitation temperature, based on SWS observations of the (0-0) S(0) and S(1) fluxes detected in the nucleus of NGC~4945, led to a cooler H$_2$ component with $T_{ex}\\sim160~\\rm K$ (corresponding to about 9\\% of the total H$_2$ mass), while a temperature $T_{ex}\\sim380~\\rm K$ (about 0.4\\% of H$_2$ mass) was estimated from the S(1) and S(2) fluxes \\citep[their Table 6]{spoon00}, which is in close agreement with our result considering the size of the big SWS slit. This excitation temperature is similar to the temperature $T_{ex} = 365\\pm50~\\rm K$ derived for the low-energy transitions in NGC~1377 \\citep{roussel06}, but higher than the temperature of $T_{ex} = 292\\pm6~\\rm K$ estimated from the S(0)-S(3) lines observed in NGC~6240 \\citep{higdon06}.  Since we cannot use the map of the H$_2$ S(3) 9.7 \\mum\\ line due to the flux mismatch between the SH and SL modules mentioned in Sec.~\\ref{sec:results}, we would only be able to base our estimate of the warm H2 mass on the S(1) and S(2) lines. In fact, even if we had been able to use the S(3) line, most of the mass is revealed by including the S(0) line. This line has however not been mapped, since the spatial resolution of the LH module is much lower than that of the SH and SL modules.            \\subsection{AGN dominated [Ne V] emission?}\\label{sec:NeV-dominated}  The detection of \\nev\\ emission from the nucleus of NGC4945 is not surprising. The galaxy hosts a 1.4$\\times$10$^6$ M$_{\\odot}$ super massive black hole, as revealed by H$_2$O maser observations \\citep{greenhill97} with an intrinsic 0.1--200\\,keV luminosity of 1.8$\\times$10$^{43}$ erg s$^{-1}$ \\citep{guainazzi00}.  More remarkable is the faintness of the \\nev\\ emission. Only 10$^{-6}$ of the infrared luminosity of the galaxy is detected in the \\nev\\ line (Table~\\ref{tab-c4:fluxes-fov}). This is surprisingly little for an AGN which can easily account for the entire bolometric emission of the galaxy \\citep{marconi00}. It is also 100 times less than L(\\nev)/L$_{\\rm IR}$ for galaxies whose IR luminosity is dominated by AGN activity \\citep{goulding09}.  Interesting is also the size of the \\nev\\ emitting region, which can be reasonably well constrained given that at a distance of 3.82\\,Mpc one arcsecond corresponds to 18.5\\,pc. Since the spatial profile of the \\nev\\ map (Figure~\\ref{fig:SH-maps1}) is only marginally wider than the PSF, the radius of the \\nev\\ emitting region will be less than 55 pc ($\\sim$3$''$). This upper limit is large in comparison to the size of coronal line regions \\citep[e.g. $\\sim$19 pc for Circinus as probed by the 2.48\\mum\\ \\SiIV\\ line;][]{prieto04}, but small compared to the kpc scale of narrow line regions as probed with \\oiii\\ \\citep[e.g. Circinus;][]{marconi94}.  Both the faintness and the observed extent of the \\nev\\ emission are likely affected by strong extinction within the nucleus. However, if the narrow line region were to extend above the plane of the molecular disk in which the AGN and the circumnuclear starburst are embedded, the amount of extinction would be less than probed by the silicate absorption feature ($A_{\\rm V}$=50-60). The amount will, however, certainly be less than the line of sight extinction to the AGN broad line region, N$_H$=2.4$\\times$10$^{24}$ cm$^{-2}$, \\citep{guainazzi00}, which is equivalent to $A_{\\rm V}$=1300.  The extinction probed by the silicate absorption feature is consistent with the F(\\nev)/F(14-195\\,keV) ratio found for Sy1 galaxies \\citep{weaver10} when multiplied by the absorption corrected 14-195\\,keV flux for NGC\\,4945. The latter depends on the geometry of the absorber, and ranges between 5 and 10 times the Swift-BAT flux of NGC\\,4945 (F(14-195\\,keV)=33$\\pm$0.8 $\\times$10$^{-18}$ W cm$^{-2}$; \\citet{tueller10}) for a spherical and edge-on disk geometry, respectively (K. Iwasawa, private communication). Taking further into account the factor 2.5 scatter in the F(\\nev)/F(14-195\\,keV) ratio for Sy1 galaxies, the attenuation on the \\nev\\ flux ranges between 12 and 157, which corresponds to $A_{\\rm V}$=55-112 mag for our adopted local ISM extinction law of \\citet{chiar06}. This result would indicate that the NLR is buried along with the other components of the nuclear molecular disk.  The strong obscuration to the NLR in NGC\\,4945 may serve as a warning that mid-infrared NLR tracers such as \\nev\\ and \\oiv\\ may not always be suitable as tracers of the AGN luminosity as suggested by \\citet{goulding09}. Galaxies in this risk group may be recognized by the presence of a deep silicate absorption feature in their mid-infrared spectra and a classification 3A, 3B or 3C in the mid-infrared galaxy classification scheme of \\citet{spoon07}. Besides NGC\\,4945 galaxies in this group include LIRGs such as NGC\\,4418 and ULIRGs like Arp220, IRAS\\,08572+3915 and IRAS\\,F00183-7111.  The example of NGC~4945 may also serve as a warning that the non-detection of NLR tracers in AGN does not necessarily imply the absence of a NLR. Instead it may indicate that the NLR is so strongly attenuated that the line does not emerge above the continuum noise. In these cases, only deep mid-IR observations, like the ones presented here, would reveal these tracers.  In NGC4945 the \\nev\\ line was detected at 5\\% of the continuum flux in a spectrum with continuum S/N of 140 in the 14-15\\mum\\ range. Other galaxies which may host an AGN, and which have been observed at similar S/N (among them ULIRGs like Arp220 and IRAS\\,08572+3915), do not show evidence for \\nev\\ emission \\citep{armus07}. The NLR in these galaxies may hence be even more strongly obscured than in NGC~4945, or the AGN may contribute a far smaller fraction of the bolometric power than in NGC~4945. The NLR may also simply not exist if the ionizing photons needed to form the NLR are 4$\\pi$ obscured close to their origin.   \\subsection{[Ne V] emission from supernovae? the template Cassiopeia A}\\label{sec:NeV-SNR}  Although the [NeV] emission is unresolved along the major axis and only marginally resolved along the minor axis (as discussed in Sec.~\\ref{sec:results}), the fact that the \\nev\\ line has also been observed in supernova remnants (SNRs) \\citep[e.g.][]{oliva99, smith09} raises valid questions: could the \\nev\\ emission observed in NGC~4945 be powered purely by SNRs?, and if so, how many SNRs would be needed to reproduce its \\nev\\ flux? Below we explore the alternative of \\nev\\ emission being powered by the SNRs in the nuclear region of NGC~4945. We note, though, that the number of supernovae needed depends not only on the total \\nev\\ flux but also on the actual size of the SNR.  Our analysis is based on previous Spitzer/IRS SL module observations of Cassiopeia A (Cas A) from which the global distribution of fine-structure lines covering a 5.3$\\times$5.3 arcminute$^2$ area was presented by \\citet{smith09}. The total \\nev\\ flux of Cas A was not reported before due to the difficulties to extract the flux with enough reliability. Using fitting techniques tailored to the \\nev\\ line, we now are able to state its total flux as $1.5\\pm0.7\\times10^{-18}~\\Wcm$. Considering an average extinction $A_{\\rm V}\\sim5$ mag \\citep{hurford96} the corrected flux, using the extinction law for the local ISM by \\citet{chiar06}, is $\\sim1.9\\pm0.9\\times10^{-18}~\\Wcm$.  On the other hand, the total \\nev\\ flux in NGC~4945 is $\\sim5.3\\pm0.5\\times10^{-21}~\\Wcm$. In order to estimate the dereddened \\nev\\ emission observed in NGC~4945, we computed an average extinction value $A_{\\rm V}\\sim36$ mag from the $A_{\\rm V}(9.85~\\mu \\rm m)$ map (Fig.~\\ref{fig:Tau-sil}). Note that due to the stronger extinction at the central pixels, this represents just a lower limit to the actual extinction on the \\nev\\ line. Thus the lower limit for the extinction-corrected total \\nev\\ flux is $\\sim$27.9$\\times10^{-21}~\\Wcm$.  Considering a canonical distance of 3.4 kpc for Cas A \\citep[][and references therein]{smith09} and $\\sim$3.82 Mpc for NGC~4945, we found that the ratio between the total \\nev\\ flux of NGC~4945 and Cas A, corrected by the square of the ratio between the distances $(D_{\\rm NGC4945}/D_{\\rm CasA})^2$, is about 19$\\times10^3$. This represents a lower limit to the number of contemporaneous Cas A type supernovae needed to produce the extinction-corrected \\nev\\ emission observed in NGC~4945. Another extreme lower limit can be obtained assuming no extinction in the NGC~4945 flux, from which the same exercise indicates that about 4$\\times10^3$ Cas A type SNRs would be needed.  According to \\citet{lenc09}, the upper limit to the supernova rate in NGC~4945 is 15.3 $\\rm yr^{-1}$, which means the individual supernovae, for the not corrected and corrected for extinction lower limits estimated above, need to emit at Cas A levels for about 260 and 1240 years, respectively. These are reasonable numbers given that the age of Cas A is estimated to be about 330 yr \\citep{fesen06}, while the kinematic age of another \\nev\\ emitting SNR, E0102, has been estimated to be $2050\\pm600$ yr \\citep{finkelstein06}. This implies that the number of SNRs estimated above can indeed persist long enough to power the \\nev\\ emission observed in NGC~4945, at least from the point of view of the SNR properties used in our estimate. However, \\citet[][their section 4.5]{lenc09} also determine a median SN age of 85 yr, which translates into a median supernova rate of $\\sim$0.12 SNe/yr, or a much (15.3/0.12$\\sim$128 times) longer radiative phase ($\\sim$3--16$\\times10^4$ years) needed for the individual supernovae. Therefore, the \\nev\\ flux observed in NGC~4945 is likely not powered by a population of supernova remnants."
    },
    "1107/1107.0012_arXiv.txt": {
        "abstract": "This is the second in a series of papers aimed at developing a practical time-domain method for self-force calculations in Kerr spacetime. The key elements of the method are (i) removal of a singular part of the perturbation field with a suitable analytic ``puncture'' based on the Detweiler--Whiting decomposition, (ii) decomposition of the perturbation equations in azimuthal ($m$-)modes, taking advantage of the axial symmetry of the Kerr background, (iii) numerical evolution of the individual $m$-modes in 2+1-dimensions with a finite difference scheme, and (iv) reconstruction of the physical self-force from the mode sum. Here we report an implementation of the method to compute the scalar-field self-force along circular equatorial geodesic orbits around a Kerr black hole. This constitutes a first time-domain computation of the self force in Kerr geometry. Our time-domain code reproduces the results of a recent frequency-domain calculation by Warburton and Barack, but has the added advantage of being readily adaptable to include the back-reaction from the self force in a self-consistent manner. In a forthcoming paper---the third in the series---we apply our method to the gravitational self-force (in the Lorenz gauge). ",
        "introduction": "}  The context of this work is the ongoing program to apply the theory of self forces (SFs) in curved spacetime to the problem of a massive compact body orbiting a black hole \\cite{Poisson:Pound:Vega:2011,Barack:2009}. This program is now reaching fruition. Recent highlights include a calculation of the gravitational SF along bound (but otherwise generic) geodesic orbits in Schwarzschild geometry \\cite{Barack:Sago:2010}, the quantitative determination of several gauge-invariant post-geodesic effects associated with the gravitational SF \\cite{Detweiler:2008,Barack:Sago:2009,Barack:Sago:2011}, and the first meaningful comparisons with post-Newtonian results \\cite{Blanchet:Detweiler:LeTiec:Whiting:2010a,Blanchet:Detweiler:LeTiec:Whiting:2010b,Barack:Damour:Sago:2010}. Results from SF calculations are being used to inform the development of an Effective One Body (EOB) model of the inspiral \\cite{Damour:2009,Barack:Damour:Sago:2010}, and even as a benchmark for fully-nonlinear numerical-relativistic simulations \\cite{NR}. Despite this significant progress, gravitational SF study so far has been restricted to the relatively simple model where the background geometry is that of a non-rotating (Schwarzschild) black hole. In order for the program to become truly astrophysically relevant \\cite{Barack:Cutler:2004}, the community must now come to address the problem of the gravitational SF in Kerr geometry.  A first calculation of the SF in Kerr spacetime has in fact been presented recently \\cite{Warburton:Barack:2010,Warburton:Barack:2011}, employing the scalar-field SF as a (relatively) simple toy model for the gravitational SF itself. This calculation---first performed for a circular orbit in the equatorial plane of the Kerr black hole \\cite{Warburton:Barack:2010} and later extended to eccentric orbits (still in the equatorial plane) \\cite{Warburton:Barack:2011}---is based on a frequency-domain treatment of the field equations: the (scalar field) perturbation equations are integrated numerically mode-by-mode in a Fourier-harmonic decomposition, and the physical scalar-field SF is then constructed using standard mode-sum regularization \\cite{Barack:Ori:2000}. This approach is computationally efficient and relatively easy to implement (as it reduces the computational task to, essentially, the solution of ordinary differential equations). However, the method also has several drawbacks: it becomes rapidly less efficient with increasing eccentricity; it cannot be incorporated easily in a scheme that evolves the orbit self-consistently under the back-reaction effect of the SF; and it is not easily generalizable to the gravitational SF in Kerr. (An alternative frequency-domain approach, designed specifically to deal with the last of the above challenges, is being developed by Shah {\\it et al}.\\ \\cite{Keidl:Shah:Friedman:2010, Shah:Keidl:Friedman:2011}.)  Time-domain methods offer a natural way around the above difficulties. Numerical codes based on time evolution are versatile and flexible, and can accommodate different orbital configurations with relatively little alteration. They can, in principle, be readily adapted to incorporate the back reaction from the SF on the orbit in a self-consistent manner. The metric perturbation equations in the Lorenz-gauge lend themselves most naturally to a time-domain treatment thanks to their hyperbolic nature. Lastly, time-domain codes formulated in 2+1D or 3+1D do not rely on separability of the perturbation equations and thus provide a natural route to the Kerr problem. For these (and other) reasons, the time-domain approach to SF calculations has been favoured by a number of workers in the field, and is under active development by several groups \\cite{Barack:Sago:2007, Vega:Diener:Tichy:Detweiler:2009,Thornburg:2009,Thornburg:2010,Canizares:Sopuerta:2009,Canizares:Sopuerta:Jaramillo:2010,Dolan:Barack:2011}.  The obvious weakness of time-domain formulations lies in their computational costliness. The lack of separability of the perturbation equations into multipole harmonics on the Kerr background means that one has to deal with partial differential equations in 2+1D (or in full 3+1D). An added complexity is introduced by the presence of a pointlike source: the physical perturbation diverges along the worldline of a pointlike particle (the divergence is Coulomb-like in the 3+1D case and logarithmic in the 2+1D case \\cite{Barack:Golbourn:2007}), and a naive numerical implementation would fail. One first has to ``regularize'' the field equations in order to make them amenable to numerical treatment. A regularization scheme formulated in 2+1D was introduced in Refs.\\ \\cite{Barack:Golbourn:2007,Barack:Golbourn:Sago:2007}, and a closely related method is being developed by Vega and collaborators \\cite{Vega:Detweiler:2008,Vega:Diener:Tichy:Detweiler:2009,Vega:Wardell:Diener:2011} within a 3+1D framework. Both methods are based on the idea of a ``puncture'' function: Rather than evolving the physical (singular) field, one evolves a ``punctured'' version thereof, obtained by subtracting from the full field a certain (singular) puncture function, given analytically. In practice, this amounts to (numerically) solving the field equations with a certain ``effective source'' to obtain the residual, regularized field. If a suitable puncture function is chosen (see below), the physical SF can be constructed directly from the residual field.  The 2+1D scheme of Refs.\\ \\cite{Barack:Golbourn:2007,Barack:Golbourn:Sago:2007} takes advantage of the axial symmetry of the Kerr background. The perturbation equations are separated into azimuthal ($\\propto e^{im\\varphi}$) modes and solved mode-by-mode, using an $m$-mode version of the puncture scheme. The SF is then reconstructed as a sum over $m$-mode contributions (see below for a more precise description). In Ref.\\ \\cite{Dolan:Barack:2011} (hereafter `Paper I') we presented a first full implementation of this {\\it $m$-mode regularization scheme}, as applied to the scalar-field SF on a Schwarzschild background. We used this simplified model as a convenient platform for development, and in order to illustrate the generic properties of the method in a relatively ``clean'' environment.  In the current work we take an important step forward by applying our method to the Kerr case. We thus present the first time-domain SF calculation in Kerr geometry. As in paper I, we consider a scalar-field toy model, and specialize to circular orbits (in the equatorial plane of the Kerr black hole). This lays the groundwork for a gravitational SF implementation in Kerr, which we will present in the third paper of the series \\cite{in_progress}.  In the remainder of this introductory section we review the $m$-mode regularization method in some detail, and give a short preview of the current work.   \\subsection{$m$-mode regularization\\label{subsec:m-mode-reg}}  Below is a schematic description of the $m$-mode regularization method as applied to the scalar-field SF in an axisymmetric spacetime such as Kerr (the basic structure of the scheme remains unaltered in the gravitational case). We model the compact object as a point particle carrying a scalar charge $q$ and moving along a geodesic of the background (e.g., Kerr) geometry. In our model, the scalar field $\\Phi$ sourced by the particle is taken to satisfy the minimally-coupled Klein Gordon equation, \\beq \\label{KG} \\Box\\Phi=S , \\eeq where the covariant derivatives in $\\Box\\equiv \\nabla^{\\alpha}\\nabla_{\\alpha}$ (as elsewhere in this work) are taken with respect to the background (e.g.\\ Kerr) metric $g_{\\alpha\\beta}$, and the source $S$ is that associated with the energy-momentum of the pointlike charge $q$ [modelled by a delta-function distribution; cf.\\ Eq.\\ (\\ref{S-3d}) below]. In our notation, $\\Phi$ represents the ``physical'', retarded solution to Eq.\\ (\\ref{KG}) (which is, of course, divergent at the particle's location). Let us recall that the physical SF acting on the scalar charge can be derived by means of the Detweiler--Whiting formula \\cite{Detweiler:Whiting:2003} \\beq \\label{Fself} F^{\\alpha}_{\\rm self}=q\\nabla^{\\alpha}\\Phi_R=q\\nabla^{\\alpha}(\\Phi-\\Phi_S). \\eeq Here $\\Phi_S$ (``S field'') is a specific singular solution of the inhomogeneous equation (\\ref{KG}), and the difference $\\Phi_R\\equiv \\Phi-\\Phi_S$ (``R field'') is a smooth, homogeneous solution of the sourceless field equation, $\\Box\\Phi_R=0$. The formal construction of $\\Phi_S$ is described in \\cite{Detweiler:Whiting:2003}.  In most cases of interest, the singular field $\\Phi_S(x)$ cannot be evaluated in exact form. However, the form of $\\Phi_S(x)$ near the particle's worldline can be studied analytically to inform a local approximation $\\Phipunc(x)$---the ``puncture function''---which may be defined globally through some extension far from the worldline. The ``order'' of the puncture describes how well it approximates $\\Phi_S(x)$ near the particle. In Paper I we introduced the following convention: a puncture is said to be of ``order $n$'', denoted $\\Phipunc^{[n]}$, if the difference $\\Phipunc^{[n]}-\\Phi_S$ is a $C^{n-2}$ function on the particle's worldline. In addition, we require $\\Phipunc^{[n]}=\\Phi_S$ as  well as $\\nabla^{\\alpha}\\Phipunc^{[n]}=\\nabla^{\\alpha}\\Phi_S$ at the particle limit (noting that the first of these conditions makes sense only for $n\\geq 2$, and the second only for $n\\geq 3$). Since the R field $\\Phi_R$ is smooth ($C^{\\infty}$), the {\\it residual} field associated with the $n$-th order puncture, \\beq \\Phires^{[n]}\\equiv \\Phi-\\Phipunc^{[n]} = \\Phi_R-(\\Phipunc^{[n]}-\\Phi_S) ,   \\label{eq-residual} \\eeq is a $C^{n-2}$ function on the worldline. Hence, for example, a 2nd-order residual function $\\Phires^{[2]}$ is continuous but not differentiable, and a 3rd-order residual function $\\Phires^{[3]}$ is both continuous and differentiable. For $n\\geq 3$ we have $\\nabla^{\\alpha}\\Phires^{[n]}=\\nabla^{\\alpha}\\Phi_R$ at the particle limit, which allows us to write \\beq\\label{Fself2} F^{\\alpha}_{\\rm self}=q\\nabla^{\\alpha}\\Phires^{[n\\geq 3]}. \\eeq  Thus, the problem of calculating the SF reduces to the problem of computing the residual field $\\Phires^{[n]}$ (for some $n\\geq 3$) near the particle. This field satisfies a ``regularized'' version of the Klein-Gordon equation: \\beq \\label{KGreg} \\Box\\Phires^{[n]} = S(x) - \\Box \\Phipunc^{[n]} \\equiv \\Seff^{[n]}. \\eeq The {\\em effective source} $\\Seff^{[n]}(x)$ is an extended source that no longer involves a delta function; it is generally $C^{n-4}$ on the worldline. Hence, a 4th-order puncture gives rise to a continuous effective source, in general. Following the development of efficient computational tools for constructing high-order Green function expansions \\cite{Ottewill:Wardell,Vega:Wardell:Diener:2011,Wardell:thesis}, a 4th-order puncture for generic orbits in Kerr geometry is now at hand, and higher-order extensions are possible.  The 3+1D treatment of Vega {\\it et al.}\\ \\cite{Vega:Detweiler:2008,Vega:Diener:Tichy:Detweiler:2009,Vega:Wardell:Diener:2011} tackles the field equation (\\ref{KGreg}) directly. In the 2+1D formulation, instead, we take advantage of the azimuthal symmetry of the background in order to reduce the dimensionality of the problem. Let $\\varphi$ be a suitable ``azimuthal'' coordinate associated with the axial symmetry of the background (below we will make a specific choice for $\\varphi$ in the Kerr case; note that the Boyer--Lindquist coordinate $\\phi$ does not suite our purpose as it is singular at the event horizon of the Kerr black hole). We formally decompose the residual field into azimuthal modes in the form \\beq \\Phires=\\sum_{m=-\\infty}^{\\infty} \\Phires^{m}e^{im\\varphi},  \\label{mmode-decomp} \\eeq and similarly for $\\Phipunc$ and $\\Seff$ (we have omitted the labels $[n]$ for brevity). This decomposition separates the field equation (\\ref{KGreg}), with each of the $m$-modes satisfying an equation of the form \\beq \\label{KGregm} \\Box^m\\Phires^{m} = \\Seff^{m}, \\eeq where $\\Box^m$ is a certain $m$-dependent differential operator in 2+1D [cf.\\ Eq.\\ (\\ref{waveeq-3}) below]. Given the modes  $\\Phires^{m}(x)$ near the particle, the SF can be reconstructed as a sum over modal contributions, \\beq \\label{FselfSum} F^{\\alpha}_{\\rm self}=\\sum_{m=-\\infty}^{\\infty} F^{\\alpha m}_{\\rm self}, \\quad\\quad F^{\\alpha m}_{\\rm self}\\equiv q\\nabla^{\\alpha}\\left(\\Phires^{m} e^{im\\varphi}\\right), \\eeq evaluated at the particle. That the gradient indeed commutes with the modal sum (for $n\\geq 2$) was shown in Ref.\\ \\cite{Barack:Golbourn:Sago:2007}. Ref.\\ \\cite{Barack:Golbourn:Sago:2007} also argued that the reconstruction formula (\\ref{FselfSum}) applies, in fact, for any $n\\geq 2$, even though the original 3+1D formula (\\ref{Fself}) makes sense only for $n\\geq 3$. Note that the individual modal contributions $F^{\\alpha m}_{\\rm self}$ depend on the puncture order $n$; however, the total SF $F^{\\alpha}_{\\rm self}$ is, of course, $n$-independent (for $n\\geq 2$).  In our method, the fields $\\Phires^{m}(x)$ are obtained by solving Eq.\\ (\\ref{KGregm}) numerically for each $m$ using a finite-difference algorithm on a 2+1D mesh. The boundary conditions for $\\Phires^{m}$ depend, in principle, on the global extension of the puncture $\\Phipunc$. It is convenient to truncate the support of $\\Phipunc$ far from the particle, in which case the residual $\\Phires$ (and its modes $\\Phires^{m}$) satisfy the usual ``retarded'' boundary conditions (outgoing radiation at infinity and purely ingoing radiation across the horizon; boundary conditions for any nonradiative modes are determined from regularity). This truncation is achieved in our scheme by introducing an auxiliary worldtube around the particle's worldline (in the 2+1D space), as we shall describe in Sec.~\\ref{sec:worldtube}.  In Paper I we implemented the above ``$m$-mode regularization scheme'' to compute the SF acting on a scalar charge in a circular orbit around a Schwarzschild black hole (refraining from a multipole decomposition and working instead in 2+1D). We implemented our numerical code with punctures of second, third, and fourth orders. In all cases our numerical results were found to be in good agreement with those of previous calculations using frequency-domain methods or time-domain integration in 1+1D, thus reaffirming the validity of the reconstruction formula (\\ref{FselfSum}) for $n=2$--$4$.  We used the simple Schwarzschild model in Paper I to explore some of the generic features of the $m$-mode method. The question of the convergence rate of the modal sum in Eq.\\ (\\ref{FselfSum}) is a crucial one, since in actual numerical implementations one can only compute a finite number of modes (realistically, no more than $\\sim 20$ modes using our current code). We made the following empirical observations (supported by heuristic arguments): (i) The {\\it dissipative} piece of the SF has an exponentially convergent mode sum; (ii) the {\\it conservative} piece of the SF has a mode-sum that converges with an inverse-power law: the individual modal contributions in the mode sum fall off at large $m$ as $m^{-n}$ for even $n$ and as $m^{-n+1}$ for odd $n$ (where $n$ is the order of the puncture used). Thus, the large-$m$ modal contribution to the conservative piece of $F^{\\alpha}_{\\rm self}$ is $\\propto m^{-2}$ with a 3rd-order puncture, and $\\propto m^{-4}$ with either a 4th or a 5th-order puncture. We concluded that a 4th-order puncture implementation offers an optimal balance between simplicity (of the analytic puncture formulation) and efficiency (of the computational algorithm). We expect this conclusion to carry over to the Kerr case.  In our implementation (as in any other time-domain implementation) the initial data for the time evolution are, of course, unknown. Instead, we begin the evolution with arbitrarily selected initial data, and rely on the property of homogeneous perturbations in black hole geometries to completely ``dissipate away'' at late time (``no hair theorem''). We let the evolution proceed until the solution relaxes to a stationary state at a sufficient level. It is important to understand the rate of relaxation, since this determines the necessary evolution time and thus has important effect on the computational burden. As discussed in Paper I, the theoretical expectation is for the modal contributions to relax with an $m$-dependent inverse-power law in time $t$. Modes with higher $m$ relax faster, and the slowest relaxation is exhibited by the $m=0$ mode: $\\propto t^{-2}$ for the field modes $\\Phires^{m}$ and $\\propto t^{-3}$ for the force modes $F^{\\alpha m}_{\\rm self}$. This behavior was confirmed numerically in Paper I; we expect it, too, to carry over to the Kerr case. The issue points to an important advantage of the 2+1D formulation (as compared with the full 3+1D treatment): Since different $m$-modes are evolved separately, we have the flexibility of being able to adjust the evolution time as a function of $m$ to achieve computational saving.   \\subsection{This paper}  Here we extend the analysis of Paper I to the Kerr case, focusing again on circular motion (in the equatorial plane). Our 2+1D treatment makes such a generalization straightforward {\\em in principle}. However, there are several nontrivial technical details that require careful consideration and extra development. First, the analytic expressions involved in the puncture formulation become considerably more complicated and require commensurably greater care in their numerical implementation. Second, the azimuthal coordinate used to define the $m$-mode decomposition must be chosen judiciously; the standard Boyer-Lindquist coordinate suits for purpose in Schwarzschild geometry but not in the Kerr case. The third point is most crucial: We find that the finite-difference scheme employed in Paper I to evolve the field equation (\\ref{KGregm}) becomes unstable in the Kerr case and must be replaced with an alternative method. We have experimented with several alternatives, and will report here our findings. We observe that some of these methods, used previously in the literature to study vacuum perturbations on Kerr, develop instabilities at large values of $m$, which makes them unsuitable for us. We develop and implement a finite-difference scheme (a variant of the ``method of lines'') that appears to perform well even for large $m$ values. Our scheme is second-order convergent (in the grid spacing), and we implement it here with a puncture of 4th order ($n=4$).  The remainder of this paper is structured as follows. In Sec.~\\ref{sec:scalar-field} we describe our physical setup and review the formalism of scalar-field perturbations on a Kerr background. In Sec.~\\ref{sec:punc_and_source} we formulate the $m$-mode regularization scheme as applied to the problem at hand, and prescribe 2nd, 3rd and 4th-order puncture functions with corresponding effective sources. Section \\ref{sec:numerical_method} describes our numerical method, including the worldtube technique and the finite-difference scheme. We also report on some of our less-successful experiments with other schemes. In Section \\ref{sec:results} we present some numerical results along with convergence tests, and explore the late-time relaxation of the numerical solutions and the convergence properties of the $m$-mode sum. We demonstrate that our code successfully reproduces the results of Warburton--Barack's frequency domain calculation \\cite{Warburton:Barack:2010}. In the concluding section, Sec.\\ \\ref{outlook}, we discuss future directions, including the application of our method to the gravitational SF.  Throughout this work we use metric signature ${-}{+}{+}{+}$ and set $G=c=1$. ",
        "conclusions": ""
    },
    "1107/1107.1708_arXiv.txt": {
        "abstract": "{X-Shooter is the first second-generation instrument for the ESO-Very Large Telescope. It as a spectrograph covering the $300-2480$ nm spectral range at once with a high resolving power. These properties enticed us to observe the well known trans-Neptunian object (136199) Eris during the science verification of the instrument. The target has numerous absorption features in the optical and near-infrared domain which has been observed by different authors, showing differences in their positions and strengths. } {Besides testing the capabilities of X-Shooter to observe minor bodies, we attempt at constraining the existence of super-volatiles, e.g., CH$_4$, CO and N$_2$, and in particular try to understand the physical-chemical state of the ices on Eris' surface. } {We observed Eris in the $300-2480$ nm range and compared the newly obtained spectra with those available in the literature. We identified several absorption features, measuring their positions and depth and compare them with those of reflectance of pure methane ice obtained from the optical constants of this ice at 30 K to study shifts in their positions and find a possible explanation for their origin. } {We identify several absorption bands in the spectrum all consistent with the presence of CH$_4$ ice. We do not identify bands related with N$_2$ or CO. We measured the central wavelengths of the bands and find variable shifts, with respect to the spectrum of pure CH$_4$ at 30 K. } {Based on these wavelength shifts we confirm the presence of a dilution of CH$_4$ in other ice on the surface of Eris and the presence of pure CH$_4$ spatially segregated. The comparison of the centers and shapes of these bands with previous works suggest that the surface is heterogeneous. The absence of the 2160 nm band of N$_2$ can be explained if the surface temperature is below 35.6 K, the transition temperature between the alpha and beta phases of this ice.  Our results, including the reanalysis of data published elsewhere, point to an heterogeneous surface on Eris. } \\titlerunning{Eris through the eyes of X-Shooter} ",
        "introduction": "To interpret the surface composition of minor bodies it is necessary to observe their spectra in the widest possible spectral range. This is typically achieved by observing the target with different instruments, techniques, and calibrations. This results in a lack of simultaneous datasets and having often to rely on other parameters (e.g., albedo, photometric magnitudes) to scale the complete spectrum. Luckily, new instrumentation allows us to overcome some of these problems. One of such new instruments is  X-Shooter (XS hereafter), which is also the first second-generation instrument developed for the ESO-Very Large Telescope (D'Odorico et al. \\cite{dodor06}) and is currently mounted in its unit 2 (Kueyen).  More detailed discussion of XS is given in Section 2, but, as an overview, XS is a spectrograph that obtains spectra from 300 to 2480 nm in a single image at high resolving powers (from 3,000 up to 10,000). This wavelength range sufficiently covers many known ices and silicates. At the time of its science verification we decided to test its capabilities by observing a well known trans-Neptunian object with absorption features in most of this spectral range: (136199) Eris (hereafter Eris).  Eris, formerly known as 2003 UB$_{313}$, is one of the largest bodies in the trans-Neptunian region, with a diameter $<2340$ km (B. Sicardy, personal communication), and, together with other two thousand-km-sized objects [(134340) Pluto and (136472) Makemake], shows CH$_4$ absorption bands in the spectrum, from the visible to the near infrared, (see Brown et al. \\cite{brown05}, Licandro et al. \\cite{lica06b}).  Rotational light-curves of Eris have had little success obtaining a reliable rotational period (e.g., Carraro et al. \\cite{carra06} or Duffard et al. \\cite{duffa08}). Roe et al. (\\cite{roe08}) proposed a rotational period of 1.08 days. The light-curve seems not related to Eris' shape, but probably to an albedo patch. Nevertheless, due to the noise in the data, even if this albedo patch indeed exists, it is clear that there are no large albedo heterogeneities.  A close look at the spectroscopic data shows that there exist a subtle heterogeneity. Abernathy et al. (\\cite{abern09}) compared their spectra with that of Licandro et al. (\\cite{lica06b}) finding different central wavelengths and different depths for a few selected methane absorption bands. Nevertheless, recent results pointed otherwise: Tegler et al. (\\cite{tegle10}), using their own results, and those of Quirico and Schmitt (\\cite{quir97a}), showed that a single component, CH$_4$ diluted in N$_2$, could explain the wavelength shifts because those shifts change by increasing with increasing wavelength.  The orbit of Eris is that of a scattered disk object (Gladman et al. \\cite{gladm08}).  It has a perihelion distance of 49.5 AU, and it is currently located at 97 AU and approaching perihelion. At 97 AU, its equilibrium temperature is $\\lesssim30$ K as estimated by the Stefan-Boltzman law and assuming conservation of energy. Therefore, based on models of volatiles retention (see Schaller and Brown \\cite{schal07}) and its similarities with Pluto, Eris' surface is expected to be covered in N$_2$, CO and CH$_4$ ices. Of these species, only CH$_4$ has been clearly seen on Eris.  Therefore, we requested observing time during XS science verification to make the most of its high resolving power to detect features of CO and N$_2$, and to obtain Eris\u2019 spectrum simultaneously between 300 and 2480 nm.  The article is organized as follows: in Sec. 2 we describe the instrument and the observations. In Sec. 3 we present the results obtained from our analysis, while the discussion and conclusions are presented in Secs. 4 and 5 respectively. ",
        "conclusions": "Using X-Shooter, a new instrument at the ESO-Very Large Telescope, we obtained two new spectra of Eris. The spectra were obtained at high resolving power, $\\lambda/\\Delta\\lambda\\sim5,000$, covering the 300 to 2480 nm at once. We compared these datasets with those available in the literature and with a dataset of Pluto.\\\\  \\noindent The main results are summarized below: \\begin{itemize}  \\item The deeper CH$_4$ absorption bands on Eris indicate a higher content of CH$_4$ than on Pluto. This interpretation is also supported by the smaller shift of the CH$_4$ features, from the positions of pure CH$_4$, when compared to Pluto.  \\item Neither N$_2$ or CO are directly detected in our spectra, whereas both have been reported for Pluto.  \\item CH$_4$ is probably diluted in another ice, likely N$_2$ and/or CO. This can be inferred from the systematic wavelength shift of the CH$_4$ absorption bands in the Eris spectra from all individual observations.  \\item We do not see major differences between the Eris spectra, indicating that there are no large heterogeneities on the regions of its surface sampled by the two observations.  \\item We do observe indications of heterogeneity at a more subtle level. Central wavelengths of individual CH$_4$ absorption bands have different shifts when independent spectra of Eris are compared. Considering only the XS data for Eris, there is an indication of an increasing blue shift with increasing wavelength. This nicely illustrates the advantage the XS design for simultaneously obtaining data over a broad wavelength range.  \\end{itemize}  \\vspace{0.5cm}"
    },
    "1107/1107.4631_arXiv.txt": {
        "abstract": "We describe a new formalism to fit the parameters $\\alpha$ and $\\beta$ that are used in the SALT2 model to determine the standard magnitudes of Type Ia supernovae.  The new formalism describes the intrinsic scatter in Type Ia supernovae by a covariance matrix in place of the single parameter normally used.  We have applied this formalism to the Sloan Digital Sky Survey Supernova Survey (SDSS-II) data and conclude that the data are best described by $\\alpha=0.135_{-.017}^{+.033}$ and $\\beta=3.19_{-0.24}^{+0.14}$, where the error is dominated by the uncertainty in the form of the intrinsic scatter matrix.   Our result depends on the introduction of a more general form for the intrinsic scatter of the distance moduli of Type Ia supernovae than is conventional, resulting in a larger value of $\\beta$ and a larger uncertainty than the conventional approach.  Although this analysis results in a larger value of $\\beta$ and a larger error, the SDSS data differ (at a 98\\% confidence level)  with $\\beta=4.1$, the value expected for extinction by the type of dust found in the Milky Way.  We have modeled the distribution of supernovae Ia in terms of their color and conclude that there is strong evidence that variation in color is a significant contributor to the scatter of supernovae Ia around their standard candle magnitude. ",
        "introduction": "Type Ia supernovae (SN) have been extensively used as ``standard candles'' to measure astronomical distances.  However, they are often called ``standardizable candles'' to emphasize the fact that the light curve shape and color can be used to correct the apparent  magnitudes, resulting in a more accurate measurement of the distance.  The correlation between light curve shape and magnitude was demonstrated by \\citet{Ph93}, who parameterized the light curve shape with the parameter ${\\Delta}m_{15}$, the decay in $B$-band magnitude 15 days after the peak.   Different parameterizations of the light curve shape parameter have been used: $\\Delta$ for MLCS2k2 \\citep{Jh07}, $x_1$ for SALT2 \\citep{Gu07,Gu10}, and $s_B$ for SiFTO \\citep{Co08}, but the parameter is essentially determined empirically to produce the tightest correlation between light curve shape and distance.  This work will use the SuperNova ANAlysis (SNANA) software package \\citep{K09a} implementation of the SALT2 model that uses $x_1$ as the light curve shape parameter.  Type Ia SN color depends on the light curve shape parameter, but SN also exhibit variations in color for fixed light curve shape.  The source of color variation is not well understood, but it is widely suspected that variations in the explosion process, host galaxy dust, and the circumstellar environment all play a role.  The difference in measured color and the mean color for a given $x_1$ is denoted by the parameter $c$.  Regardless of the process that produces different colors, it may be expected that the apparent $B$-band magnitude ($m_B$) will show some dependence on color, and that an appropriate color correction should improve the precision of SN as standard candles.  These considerations have led to the widespread practice of fitting SN light curves  with a one parameter family of templates plus a color law to describe the color.  The determination of color is straight-forward given multi-band photometric measurements and a color law, and the result is traditionally given as $E(B-V)$.  However, the amount of correction to apply to a given SN has been less clear.  The correction has been often expressed as an extinction correction, which is related to the color through the parameter $R_V=A_V/E(B-V)$, where $A_V$ is the extinction in $V$-band.  Many efforts to determine this correction have relied on the parameterization of Cardelli, Clayton, and Mathis \\citep{CCM}, which we will call the CCM model, where the parameter $R_V$ changes both the shape of the extinction curve as a function of wavelength and the amount of extinction for a given color excess $E(B-V)$.  The simplest approach \\citep{Jh07} is to fix $R_V=3.1$, the average value determined from the studies of extinction in the Milky Way \\citep{Sn78}.  However, many studies of SN light curve data have indicated that smaller values of $R_V$ provide a better description of the SN data.  In the CCM model the single parameter $R_V$ can be measured by estimating $A_V$ or by comparing the relative extinction in different filters (by considering $E(B-V)-E(V-R)$, for example).  The SALT2 model parameterizes the SN rest-frame spectrum $S(\\lambda,p)$ as a function of wavelength ($\\lambda$) and phase ($p$) according to \\begin{equation} S(\\lambda,p)=x_0 \\, [S_0(\\lambda,p)+x_1 S_1(\\lambda,p)] \\, e^{[-cCL(\\lambda)]},\\label{eqn:salt2} \\end{equation} where the fitted parameters $x_0$, $x_1$, and $c$ are the overall scale, the light curve shape parameter, and the color respectively.  The functions $S_0$ and $S_1$ are fixed by a training sample of SN light curves and spectra.  The SALT2 light curve fit does not rely on the CCM parameterization, but uses a color law $CL(\\lambda)$ determined empirically from SN data.  The SALT2 parameters are determined by fitting the measured light curve data for each SN to the synthetic magnitudes calculated from the redshifted spectrum given by Equation \\eqref{eqn:salt2}.  The light curve parameters $x_0$, $x_1$, and $c$ specify a model SN spectrum, but an additional step is needed to use the light curve parameters to calculate a standard SN magnitude. One can calculate the apparent magnitude $m_B$, where $m_B$ is the apparent $B$-band synthesized from the rest-frame spectrum given by Equation \\eqref{eqn:salt2}, and a standardized SN brightness $m_B^s$ is computed with parameters $\\alpha$ and $\\beta$: \\begin{equation} m_B^s=m_B + \\alpha x_1 - \\beta c.\\label{eqn:basic} \\end{equation} \\citet{Gu07} describe a technique that simultaneously determines the cosmological parameters and the coefficients ($\\alpha$ and $\\beta$) that minimize the residuals on the Hubble diagram.  The SALT2 approach to correction for color is formally equivalent to the description in terms of $R_V$ and $E(B-V)$, except for the interpretation that the value of $\\beta$ may result from some effect other than extinction by host galaxy dust.  Regardless of the interpretation, $\\beta$, which gives the magnitude correction for $B$-band, and $R_V$, which gives the magnitude correction for $V$-band, may be compared via the relation $\\beta=R_B=R_V+1$.  In this work, we present a new method for fitting the parameters $\\alpha$ and $\\beta$ from the data.  The method has been implemented in the program SALT2mu, which has been added to the SNANA software package.  The motivations for a new program are: \\begin{itemize} \\item The fit for $\\alpha$ and $\\beta$ is decoupled from the cosmology.  If the mean values of $x_1$ and $c$ are a function of redshift (a magnitude limited survey will have such selection effects), the average residual depends on both the cosmology and the $\\alpha$ and $\\beta$ parameters.  We prefer to determine the $\\alpha$ and $\\beta$ solely by minimizing the rms of the distribution and to let the mean of the distribution determine the cosmology. \\item A new description of intrinsic scatter is used as discussed below. \\item The SALT2mu output can be used interchangeably with different cosmology fitters since SALT2mu produces distance moduli separately from the cosmology fit.  This feature allows a direct comparison of SN distance moduli without reference to a particular cosmology. \\item The sample used to determine $\\alpha$ and $\\beta$ can be different from the sample used to determine the cosmology.  \\end{itemize}  A significant additional benefit to writing a new program is that SALT2mu reads files from the SALT2 light curve fits and outputs files that can be read by any of the SNANA cosmology fitters.  The fitting model resides in a single function making it easy to modify the $\\chi^2$ function by adding new parameters, for example.  The inclusion of a more general description of the intrinsic scatter of the fitted SALT2 parameters is a major feature, as is the determination of $\\alpha$ and $\\beta$ in a way that is independent of the cosmology.   The details of the incorporation of the intrinsic scatter are given in the mathematical description (\\S\\ref{sec:math}) below, and most of the paper is devoted to the issues associated with the more general description of the intrinsic scatter.  The program is then applied to the Sloan Digital Sky Survey supernova (SDSS-II SN) data, which is used to illustrate some of the features of the program and issues with its use.  The SDSS-II SN data \\citep{Fr08} were obtained during 2004-2007 as part of extension of the original SDSS \\citep{Yo00}.  The SDSS telescope \\citep{Gu06} and imaging camera \\citep{Gu98} produce photometric measurements in each of the 5 SDSS filters \\citep{Fu96} spanning the range of 350 to 1000 nm.  The most useful filters for observing SDSS SN, however, are \\textit{g}, \\textit{r}, and \\textit{i} because the SN are too faint in \\textit{u} and \\textit{z} to be well measured by the SDSS SN survey except at low redshifts.  Further details of the survey, data processing and selections are given in \\S\\ref{sec:SDSS}. ",
        "conclusions": "The histograms in Figure \\ref{fig:abmarg} show the same trend that was observed in Table \\ref{tab:meas}:  when the intrinsic scatter in color is large, the best fit value of $\\beta$ will increase while $\\alpha$ decreases slightly.  The fact that our value for $\\beta$ is the highest value of those reported in \\S\\ref{sec:SDSS} depends on our incorporation of a large intrinsic scatter in supernova color.  We have not made any correction for the selection bias of our sample based on the results of the simulation, which is subject to the same selection criteria, and shows no significant bias in the determination of $\\alpha$ and $\\beta$.  The error on our result is dominated by an uncertainty in the intrinsic scatter matrix, but there are  other potential problems that have not been taken into account.  Certainly, our assumption that the intrinsic scatter is a constant, independent of any SN properties, is suspect given our lack of knowledge of the processes involved.  In addition, our model assumes that the data follow the relationship of Equation \\eqref{eqn:stmag} with a fixed value of $\\beta$.  However, we know that the value of $R_V$ varies significantly in our own galaxy \\citep{Sc10}, and we should expect to see at least some variation in SN host galaxies.  Variations in $\\beta$ produce effects similar to the intrinsic scatter in color, but the effect can be magnified in a magnitude limited sample like SDSS where there are few highly reddened SN, which have higher weight in determining $\\beta$.  Variations in $\\beta$ could explain the small differences seen between the values of $\\beta$ in Figure \\ref{fig:abred} in the two lowest redshift bins.  We have seen that neglecting intrinsic scatter in the SALT2 $c$ parameter results in a value of $\\beta$ that is biased low.  The amount of bias depends on the amount of intrinsic scatter relative to the range in the variable $c$.  Since the range in $c$ decreases as the redshift increases, neglecting or underestimating the amount of intrinsic scatter in $c$ would be expected to result in a decreasing value of $\\beta$ as a function of redshift.  The fits to the SDSS data as a function of redshift include intrinsic scatter in color so this effect should not be the explanation for the SDSS data---unless the scatter was underestimated.  However, an overestimate of the intrinsic scatter in $m_x$ leads to similar biases and is more likely explanation for the lower value of $\\beta$ that is observed in the highest redshift bin.  A more problematic type of variation in $\\beta$ could arise from the color smearing.  We have argued that color variations of SN must have some component of host galaxy extinction and also some intrinsic color variation, which could be due to variations in the explosion process or local environment or both.  We have modeled the intrinsic color variation as color smearing with an effective value $\\beta=0$, but it is possible that each mechanism that contributes to the observed color requires a different correction to the distance modulus.  Our model can only produce an effective value of $\\beta$ that could be a function of color, redshift, host galaxy type or other parameters and could, therefore, differ depending on the characteristics of the particular SN sample.  Although we have not included correlations with galaxy types or spectral line features, the introduction of additional dependent variables into Equation \\eqref{eqn:stmag} is straightforward.  In fact, having a more complete description of the supernova explosion will not only provide a smaller dispersion for the numerator in Equation \\eqref{eqn:s2mchi2}, it will decrease the size of the intrinsic scatter that is required and decrease the importance of the uncertainty as to its form.  Recently a number of papers \\citep{Ki11,La11,Ma11} have explored the technique of determining the intrinsic scatter from the data.    The approach of \\citep{Ma11} is similar to ours in that it included our more general form for the intrinsic error matrix but assumed that the off-diagonal terms involving $M_B$ were zero.  We have not considered including the intrinsic scattering error as a fit parameter partly for simplicity and partly because the underlying physics of the intrinsic scatter is still uncertain.  However, the more rigorous statistical approach described by these authors is a possible direction for future work.  We argue from the SDSS SN photometric data that there is likely some intrinsic scatter in SN color that is responsible for some of the scatter in the Hubble diagram. A similar conclusion is reached by \\citet{Fo11}, based on much more detailed information involving differences in the line velocities measured from the SN spectra, which are modeled as arising from differences in viewing angle.  While this type of detailed spectral information may not be available for all future SN surveys, an understanding of the underlying physical processes will be important for the most accurate treatment of the data."
    },
    "1107/1107.3157_arXiv.txt": {
        "abstract": "{ We consider general theories of a massive spin-2 particle $\\hmn$ on a Minkowski background. A decomposition of $\\hmn$ in terms of helicity eigenstates allows us to directly test whether any given theory possesses a consistent description as a massive spin-2 representation of the Poincar\\'e group. We demonstrate (i) that any nonlinear theory with an Einsteinian derivative structure either contains ghosts or does not describe a weakly coupled spin-2 and (ii) that there exists a two-parameter family of non-Einsteinian cubic self-interactions which constitute a ghost-free massive spin-2 theory. }  \\subheader{LMU-ASC 30/11}  \\begin{document} ",
        "introduction": "The search for viable theories of massive gravity has proven to be very difficult. Massive gravity as the theory of an interaction mediated by a massive spin-2 field must necessarily propagate five degrees of freedom (DOF), namely its five polarizations. While this fixes the mass term uniquely on the linear level to the Fierz-Pauli form \\cite{fierz_pauli_1939}, it also introduces immediate troubles. When taking the massless limit of the theory, the additional graviton polarizations do not generically decouple. As a consequence, predictions of any massive gravity theory in a first linear approximation differ from massless general relativity (GR) by numerical factors \\cite{vandam_veltman_zakharov_1970}. However, it has been pointed out in \\cite{vainshtein_1972} that this so-called van Dam-Veltman-Zakharov (vDVZ) discontinuity may disappear when correctly taking nonlinearities into account, with nonlinearities growing with decreasing graviton mass as $m^{-4}$. Later, in \\cite{deffayet_etal_2001} this nonperturbatively continuous behavior has been demonstrated for a specific model of massive gravity.  The question of instabilities in nonlinear theories of massive gravity has been addressed in many works. Boulware and Deser \\cite{boulware_deser_1973} proved the inevitable appearance of the sixth polarization of the graviton as a ghost-like state in a wide class of models through a Hamiltonian formalism. On the level of the action, the appearance of a six-derivative cubic term in the helicity-0 component in Einstein gravity with a Fierz-Pauli mass term has been shown in \\cite{deffayet_etal_2001} in terms of the leading singularity of the graviton vertex, and in terms of St\\\"uckelberg fields in \\cite{arkanihamed_etal_2002}. Both in \\cite{arkanihamed_etal_2002, creminelli_etal_2005} it was proven that one may cancel such operators by appropriately adding non-derivative interactions to the action. However, it remained unsettled if other operators could spoil the stability of the theory. Since then, several other works have tried to construct manifestly stable theories of a single massive graviton (\\cite{derham_gabadadze_etal,Chamseddine:2010ub} and references therein). Recently, \\cite{derham_etal_2010} suggested that one may formulate nonlinear theories containing the Fierz-Pauli mass term, which, on the full nonlinear level, were found to describe the correct number of degrees of freedom in \\cite{HassanRosen}. Others have argued that generic theories of massive gravity contain problems such as ghosts or superluminality \\cite{alberte_etal_2010,gruzinov_2011}. While the latter is beyond the scope of this work, we will confirm findings on the ghost problem in massive gravity.  We will demonstrate that an expansion of nonlinear Fierz-Pauli models in terms of a weakly coupled massive spin-2 field inevitably leads to inconsistencies. We consider theories which are Poincar\\'{e} invariant, local and weakly coupled in at least some energy interval where they describe a single propagating massive spin-2 particle, $\\hmn$, on a Minkowski background. They can be expressed in terms of polynomial interactions of a local Lagrangian. For a stability analysis, it is important to realize that ghost-type instabilities are relevant on arbitrarily short timescales: While tachyonic instabilities arise for momenta lower than the tachyonic mass $m_\\text{tach}$ and contain an intrinsic timescale $t = 1/m_\\text{tach}$, ghosts are UV instabilities. Their intrinsic timescale is $t = 1/E$, which is only limited by an effective field theory (EFT) cutoff $\\Lambda$. It is therefore justified to test the system for instabilities at momenta $m \\ll E \\ll \\Lambda$. This gives a straightforward prescription for a stability analysis in terms of representations of the Poincar\\'{e} group. At very high energies, $E \\gg m$, the massive spin-2 representation of the Poincar\\'e group decomposes into the direct sum of irreducible helicity representations. In other words, the ratio of mass and energy $m/E$ parametrizes the mixing between helicity eigenstates. For sufficiently large $E$, this mixing becomes negligible, helicity eigenstates decouple on a linear level and constitute a unique description of the system. As we only consider weakly coupled theories, we can capture all relevant physics by decomposing $\\hmn$ \\emph{linearly} into its helicities. Any inconsistencies found in these states inevitably indicate (i) ghost-like instabilities or (ii) a violation of the assumptions of Poincar\\'e invariance, locality and weak coupling. Either way, a consistent description in terms of a weakly coupled massive spin-2 field $\\hmn$ is excluded. In this sense, decomposing the massive spin-2 into its helicities has a big advantage compared to previous methods. It enables one to not only count degrees of freedom, but gives a direct test whether they can be grouped into a massive spin-2 particle.  This work will therefore elucidate aspects of the ongoing debate by explicitly showing that no nonlinear extension of Fierz-Pauli theory with an Einsteinian derivative structure is free of inconsistencies. Two interpretations are possible. The first is that the theory contains ghosts, which must be addressed in any analysis of viability. One has to introduce additional degrees of freedom that cure the ultraviolet (UV) regime, or ensure that the ghost degrees of freedom are otherwise shielded. Note here that we are not interested in the position of ghost poles, but are simply looking for the mere appearance of pathological degrees of freedom. The second interpretation, possibly applicable to \\cite{HassanRosen}, is that the weakly coupled degrees of freedom cannot be attributed to a massive spin-2 Poincar\\'e representation or the theory exhibits strong coupling already in the infrared.  However, an Einsteinian derivative structure is not preferred when considering general theories of massive spin-2 fields. For $m=0$ it is known that Poincar\\'{e} invariance, locality and unitarity alone pin down general relativity as the unique theory of self-interactions \\cite{weinberg_1964_1965, Deser}. For $m \\neq 0$ these arguments cannot be generalized. Henceforth, we will construct a ghost-free cubic theory of a massive spin-2 field by only requiring the self-interactions to be Lorentz-invariant and to involve at most two derivatives. A priori, all coupling parameters are arbitrary and can be adjusted model dependently. We will exploit this freedom in parameters to eliminate possible ghost-like instabilities and extra DOFs and prove that the constructed action is a valid theory of a weakly coupled massive spin-2 particle.  Our paper is organized as follows. In section \\ref{sec:FP}, we review the theory of a free massive spin-2 particle due to Fierz and Pauli \\cite{fierz_pauli_1939} and explain how it can be understood in terms of helicity degrees of freedom, whose mixing is governed by the mass of the particle. We will further explain why a decomposition into helicities as kinetic eigenstates is necessarily linear.  We apply the helicity decomposition to the class of theories whose derivative interactions are governed by Einsteinian vertices in section \\ref{sec:ECI}. We demonstrate that inconsistencies inevitably appear on the cubic level, even when allowing for additional arbitrary nonderivative interactions. Furthermore we discuss recent claims \\cite{derham_gabadadze_etal, HassanRosen} in the literature regarding the eligibility of such a theory.  Finally, in section \\ref{sec:GCI}, we construct a two-parameter family of cubic interactions which manifestly describe the ghost-free theory of a self-interacting massive spin-2 particle. This two-parameter family does not contain the cubic part of the Einstein-Hilbert action and is therefore not a likely candidate for a massive version of gravity. ",
        "conclusions": "The aim of this work was to investigate the consistency of nonlinear extensions of Fierz-Pauli theory on a Minkowski background. In contrast to previous works \\cite{arkanihamed_etal_2002, derham_gabadadze_etal}, no attention was paid to restoration of general covariance. Instead we made direct use of the isometries of the flat background and worked in terms of irreducible representations of the Poincar\\'e group. In addition, the only requirements were locality and weak coupling, s.t. the theory could be understood as polynomial interactions of the field $\\hmn$. Any theory meeting these requirements is inevitably subject to our results.  Our set of prerequisites allowed us to greatly reduce the complexity of the analysis. A weakly coupled theory can, for sufficiently weak fields, be understood in terms of its lowest order interaction, which enabled us to focus solely on the cubic action. Further, we made use of the fact that for high momenta, the massive spin-2 representation decomposes into a direct sum of helicity-2, -1 and -0 representations and all physical information is contained in helicity eigenstates. Finally, weak coupling, amongst other things, requires the decomposition to be linear, as this excludes mixing between different order interactions.  We first applied this formalism to a massive gravity theory with derivative interactions governed by the expanded Ricci scalar. Despite allowing the addition of nonderivative interactions, we found higher derivative operators already on the cubic level. Furthermore, we demonstrated that in the same class of theories, the scale $\\Lambda_5^5 \\equiv m^4 M_P$ cannot be fully eliminated from the action when taking quartic interactions into account.  Our result only allows for two conclusions. Either the theory violates the prerequisites, i.e. the DOFs are not a massive spin-2 particle, they are not weakly coupled or the theory is nonlocal, or there are additional ghost degrees of freedom present. The former interpretation in particular applies to the class of models introduced in \\cite{HassanRosen}, since the authors were able to prove the absence of an additional sixth degree of freedom on the full nonlinear action. It is an interesting question what these theories really describe. First hints may be drawn from the fact that a nondynamical auxiliary metric is necessarily introduced into the action. Allowing for dynamics of this auxiliary field, albeit introducing additional degrees of freedom, may in the end provide valuable insight into this class of models.  Finally, we applied our formalism to the most general cubic Lagrangian for the field $\\hmn$. We showed that one can tune all interactions in such a way that a manifestly ghost-free two-parameter family of theories of a massive spin-2 field is found. This could be seen both in terms of helicities and in terms of components of $\\hmn$. A Dirac constraint analysis revealed the same set of conditions as in linear Fierz-Pauli theory. A possible phenomenological application of this action has to be analyzed in more detail. Since the derivative structure differs from the Einsteinian cubic vertex, it is highly doubtful that it can be applied to the problem of giving a mass to the graviton. Furthermore, it still encounters the problem of tachyonic instabilities, as any cubic theory. One may however easily extend our analysis to higher order interactions and this way be able construct manifestly stable theories of a self-interacting massive spin-2 field."
    },
    "1107/1107.4541_arXiv.txt": {
        "abstract": "{The growing interest in the high-z universe, where strongly obscured objects are present, has determined an effort to improve the simulations of dust formation and evolution in galaxies. Three main basic ingredients enter the problem influencing the total dust budget and the kind of mixture of the dust grains: the types and amounts of dust injected by AGB stars and SN{\\ae} and the accretion and destruction processes of dust in the interstellar medium (ISM). They govern the relative abundances of the gas and dust components of the ISM of a galaxy.}{In this study, we focus on the dust emitted by stars  and present a database of condensation efficiencies for the refractory elements C, O, Mg, Si, S, Ca and Fe in AGB stars and SN{\\ae} that can be easily applied to the traditional gaseous ejecta, in order to determine the amount and kind of refractory elements locally embedded into dust and injected into the ISM.} {The best theoretical recipes available nowadays in literature to estimate the amount of dust produced by SN{\\ae} and AGB stars have been discussed and for SN{\\ae} compared to the observations to get clues on the problem. The condensation efficiencies have been then analyzed in the context of a classical chemical model of dust formation and evolution in the  Solar Neighbourhood and Galactic Disk.} {Tables of condensation coefficients are presented for (i) AGB stars at varying the metallicity and  (ii) SN{\\ae} at varying the density n$_{\\textrm{H}}$ of the ISM where the SN{a} explosions took place. In particular, we show how the controversial CNT approximation widely adopted  to form dust  in SN{\\ae}, still gives good results and agrees with some clues coming from the observations. A new generation of  dust formation models in SN{\\ae} is however required to solve some contradictions that have recently emerged.} {A simple database of condensation efficiencies is set up  to be used in chemical models including the effect of dusts of different type and meant to simulate real galaxies of different type going from primordial proto-galaxies  to those currently seen in the local universe.} ",
        "introduction": "Understanding and modelling the interstellar dust has recently received a great deal of attention thanks to current observations unveiling the existence of a high-z universe heavily obscured by large quantities of dust \\citep{Shapley01,Carilli01,Bertoldi03,Robson04,Wang08a,Wang08b,Michalowski10a,Michalowski10b}. Once established the presence of dust, some key questions must be answered about the physical nature of a  dust-rich  universe. What is the origin of these copious amounts of dust? What is the dust composition? What is a plausible  mixture of the dust grains  able to account for the observational properties of extinction of the stellar light and emission in the mid and far infrared (MIR/FIR)? Starting from the first simplified models simulating in some way the formation and evolution of dust in galaxies \\citep{Lisenfeld98,Morgan03,Inoue03}, over the years models of growing complexity have been presented: from the pioneer work of \\citet{Dwek98} on the MW till the recent ones by \\citet{Zhukovska08} and \\citet{Piovan11a} on the Solar Neighborhood (SoNe) of the Milky Way (MW) or the whole Galactic Disk of \\citet{Piovan11b}, on galaxies of different morphological types \\citep{Calura08}, on star-burst galaxies \\citep{Gall11a}, on QSOs, LBGs and the Early Universe \\citep{Valiante09,Valiante11,Pipino11,Dwek11,Mattsson11,Gall11b,Yamasawa11}.\\\\ \\indent The concept of duty cycle for  the dust must be introduced to suitably describe the formation and evolution of dust in high-z and local galaxies, and to simultaneously infer  precious clues on when and how galaxies formed and evolved. The cyclic history of the interstellar dust is described in detail by \\citet{Zhukovska08} and nicely illustrated in the classical diagram by  \\citet{Jones04}. In brief,  low and intermediate AGB stars thanks to mass loss by stellar winds, and massive stars thanks to the  Core Collapse SN{a} explosions (CCSN{\\ae}), inject refractory elements in the ISM: most of this material is in gaseous form, but important amounts of it condense into the so-called star-dust. Once mixed in the turbulent ISM, star-dust grains are subjected to destructive processes that restitute the material to the gaseous phase. The competition between this process and the one of dust  accretion onto the so-called seeds in dense and cold molecular clouds (MCs), determines the total budget of dust in the ISM and the observed depletion of the refractory elements by formation of new dust grains. In the MCs, where dust accretes and cools down  the region, star formation takes place generating new  stars that in turn evolve and die, thus more and more enriching the ISM with new metals and star-dust (the fraction of it able to survive to  local shocks). It is soon evident even from this simple description  that some key agents must intervene  to drive the evolution of dust. They are identified with some grains or  grain families with given composition and properties, the physical mechanisms of formation/accretion and destruction of dust in  the ISM,  and finally the yields of  dust by stars.\\\\ \\indent In this study we focus the attention on the dust emitted by stars of different mass and metallicities during the AGB evolutionary phase and/or the SN{a} explosion as appropriated to their initial mass \\citep{Zhukovska08}. The amount of produced star-dust and the injection timescales are the object of a vivid debate, largely motivated by the high-z galaxies. Indeed, it is not clear (i) whether star-dust of SN{a} origin  is able alone to explain the amount of dust observed in high-z objects \\citep{Gall11b,Gall11b}, (ii) up to  which redshift and how strong is the role played by massive AGB stars \\citep{Valiante09,Dwek11}, and (iii) whether the contribution by the dust accreted in the ISM cannot be neglected \\citep{Dwek09,Draine09,Mattsson11}. In this study, we intend to thoroughly discuss what could be the best compilation of theoretical condensation efficiencies currently available in literature and how much of each refractory element could locally condense in form of star-dust. To this aim we present here a easy-to-use compilation of condensation efficiencies $\\delta$ \\citep{Dwek98,Calura08,Piovan11b} to be applied to the masses of \\textit{single elements} restituted by stars to the ISM. In other words, starting from the classical compilations of the gas mass in form of a given element ejected by each  star during its life back into the ISM  \\citep{Portinari98,vandenHoek97,Francois04}, we provide a compilation of coefficients giving the dust-to-gas ratio for that specific ejecta.  \\renewcommand{\\arraystretch}{1.1} \\begin{table*} \\label{DeltaLiterature} \\footnotesize \\begin{center} \\caption[]{Prescriptions taken from  literature to model the star-dust contribution to the ISM:} \\begin{tabular}{ccc}          \\hline \\hline \\noalign{\\smallskip} Work & AGB stars & SN{\\ae} \\\\   \\hline \\noalign{\\smallskip} \\citet{Calura08}$^{1}$ & \\citet{Dwek98} & \\citet{Dwek98} \\\\ \\citet{Zhukovska08}$^{2}$ & \\citet{Ferrarotti06}    &  Its own $\\delta^{SN}$ scheme     \\\\ \\citet{Valiante09}$^{3}$ & \\citet{Ferrarotti06}   &   \\citet{Bianchi07} \\\\ \\citet{Pipino11}$^{4}$ & \\citet{Dwek98} & revised \\citet{Dwek98} \\\\ \\citet{Yamasawa11}$^{5}$ & no AGB stars & \\citet{Nozawa03,Nozawa07} \\\\ \\citet{Gall11a,Gall11b}$^{6}$  & \\citet{Ferrarotti06}      & \\citet{Todini01,Nozawa03,Nozawa06}     \\\\ \\citet{Dwek11}$^{7}$  & \\citet{Dwek98}      & Its own $\\delta^{SN}$ scheme     \\\\ \\hline \\end{tabular} \\end{center} \\begin{flushleft}\\footnotesize \\footnotesize$^{1}${The same condensation efficiencies $\\delta$ proposed by \\citet{Dwek98} are adopted.}\\, \\footnotesize$^{2}${Low and constant  condensation efficiencies are proposed and adopted for SN{\\ae}.}\\, \\footnotesize$^{3}${The original model by \\citet{Todini01} for dust formation in SN{\\ae} is extended to a wider set of initial conditions and model assumptions.}\\, \\footnotesize$^{4}${Condensation efficiencies $\\delta$ by \\citet{Dwek98} are lowered to match the observation of dust in CCSN{\\ae}, thus including in some way the uncertainties of the destructive reverse shock effects.}\\, \\footnotesize$^{5}${Only SN{\\ae} are included as dust factories because the study is limited to the very early universe.}\\, \\footnotesize$^{6}${Average coefficients are obtained in order to study the evolution of the total dust mass in star-bursters and QSOs.}\\, \\footnotesize$^{7}${Only the average total amount of dust formed in SN{\\ae} and WR is considered.}\\, \\end{flushleft} \\end{table*} \\renewcommand{\\arraystretch}{1}  Many different recipes are proposed to deal with the two main factories of star-dust (AGB stars and SN{\\ae}), each of which  with a different level of complexity  \\citep{Gail09,Dwek05}: some of them consider  only the total amount of dust that is injected and neglect its composition (spectrum of elements), others adopt simple schemes to follow the evolution of a group of elements and/or molecules taken as representative of the dust in the ISM. In Table \\ref{DeltaLiterature} we list all the prescriptions we have adopted based on the most recent models of dust formation and destruction.\\\\  \\indent  The plan of the paper is as follows. In Sect. \\ref{SNEyield_details} we discuss the amount of dust injected by a SNa in the ISM, both from theoretical and observational point of view, look at the different types of SN{\\ae} producing dust, calculate and present the condensation efficiencies for the single elements from various sources in literature, and finally analyze the various alternatives highlighting their merits and drawbacks. In Sect. \\ref{AGByield} we examine the production of dust by  AGB stars. In Sect. \\ref{Yieldseffect} we analyze the prescription  for stardust we have just derived and their effects with the aid of the classical model for the Galactic Disk and SoNe of the MW by \\citet{Piovan11b}. Although the model includes the injection of star-dust, dust accretion and destruction in the ISM, radial flows of matter and effects of the Bar for the innermost regions, we limit ourselves here to examine only the effects brought about by type II SN{\\ae}, type Ia SN{\\ae}, and AGB stars in  three regions of the MW disk: an inner region, the SoNe, and an outer region. In Sect. \\ref{Discus_Concl} we summarize the results and draw some conclusions. This paper is the first of series of three \\citep{Piovan11b,Piovan11c} dedicated to the wide subject of dust formation/ destruction  and evolution, and its effects. Particular attention is paid to the MW which is the ideal workbench for any model of chemical evolution. In \\citet{Piovan11b} we will present our chemical model for the MW-SoNe with dust and formation and evolution based upon the classical model with infall developed long ago by \\citet{Chiosi80} and ever since used by many authors. In \\citet{Piovan11c} we will apply the same model to investigate the radial chemical properties of the MW Disk. ",
        "conclusions": "In this paper we have described the database of condensation efficiencies for the main elements that form the dust emitted by the most important dust factories in nature, i.e. SN{\\ae} and AGB stars. The results are organized in tables containing the condensation efficiencies for the refractory elements C, O, Mg, Si, S, Ca (for this latter and partially for S, the average values between the condensation efficiencies of the other refractory elements has been adopted) and Fe in AGB stars and SN{\\ae}. The condensation efficiencies, multiplied by the gaseous ejecta provide an estimate of the amount of each element trapped  into dust that is injected into the ISM. Our compilation stands upon the \\textit{theoretical} work by \\citet{Ferrarotti06,Zhukovska08} and \\citet{Nozawa03,Nozawa06,Nozawa07} and allows us  to take into account: (1) the different metallicity of the AGB stars, thus simulating the growing number of C-rich stars forming carbonaceous grains at lowering the metallicity; (2) the different density of the environment in which the SNa explosion occurs which crucially determines the final amount of dust surviving to the shocks. \\\\ \\indent There is an important aspect of the results to note. The mixture of dust grains emitted by stars, in particular the one by SN{\\ae}, is still very model dependent: for instance the kinetic models by \\citet{Cherchneff10}  not only predict different total amounts of dust as compared to the classical CNT models by \\citet{Nozawa03}, but also different  composition for the dust. To somehow cure this point of uncertainty we calculated (and tested) \\textit{average} condensation coefficients for \\textit{each elements}. To this aim, we made use of the best prescriptions for the production of dust by stars, and compared the results one would obtain by introducing them in the classical chemical model for the MW Disk and SoNe \\citep{Piovan11b,Piovan11c}. As already mentioned, the largest uncertainty is due to SN{\\ae}: to highlight the point we compare the model results with the available data about the amounts of dust and its composition, taking into account the observational hints by pre-solar grains and IR emission from SNa remnants. In particular we look for the best solution to be adopted for SN{\\ae}. We find that the CNT approximation widely adopted for dust formation in SN{\\ae}, despite its limitations as a first order approximation of a complex situation neither at equilibrium nor at steady state, still gives good results. Indeed, (i) it produces amounts of carbonaceous grains that agree with the estimate required by the pre-solar dust grains; (ii) it produces significant amounts of silicates that seem to be necessary to reproduce very recent cold dust emission spectra in the SN 1987A; (iii) the amount of dust produced by SN{\\ae} is not negligible, as indicated by the most recent in observations at long wavelengths and, (iv) the CNT unmixed model agrees well with both the FIR/sub-mm estimates of newly formed dust in SN{\\ae} and the amount of dust required to explain high-z obscured objects with the SN{\\ae} playing a decisive role.\\\\ \\indent The whole picture we have just outlined could be considered as satisfactory.  However, it is worth noticing: (i) the reduction in the amount of dust predicted by the more physically correct kinetic models, goes in the opposite direction with respect to what indicated by observations of cold dust; (ii) the poor data available on pre-solar silicates does not support an efficient condensation of silicates as indicated  by estimates of cold dust  in SN{\\ae}.\\\\ \\indent This simply means that our understanding of dust formation in SN{\\ae} is still very uncertain, both theoretically and observationally. More data on  cold dust for a wider sample of SN{\\ae} is needed. A new generation of models of dust formation in SN{\\ae} is therefore required to highlight the many points of uncertainty and contradiction. They go from the obvious uncertainty in the correct approach to use, the complicated description of the dust formation and the the limited number of accurate models in literature to poor statistics, lack of high quality observations of the cold dust in the FIR/sub-mm for many SN{\\ae}, and  uncertainties on estimating the amount of fresh dust).\\\\  \\textit{Acknowledgements}. L. Piovan acknowledges A. Weiss and the Max Planck Institut Fur AstroPhysik (Garching - Germany) for the very warm and friendly hospitality and providing unlimited computational support during the visits as EARA fellow when a significant part of this study has been carried out. The authors are also deeply grateful to  S. Zhukovska and H. P. Gail for many explanations and clarifications about their  yields of dust from AGB stars,  T. Nozawa and H. Umeda for many fruitful discussions about SNa dust yields. This work has been financed by the University of Padua with the dedicated fellowship \"Numerical Simulations of galaxies (dynamical, chemical and spectrophotometric models), strategies of parallelization in dynamical lagrangian approach, communication cell-to-cell into hierarchical tree codes, algorithms and optimization techniques\" as  part of the AACSE Strategic Research Project.  \\appendix"
    },
    "1107/1107.2628_arXiv.txt": {
        "abstract": "We investigate theoretical and observational aspects of  a time-dependent parameterization for the dark energy equation of state (EoS) $w(z)$, which is a well behaved function of the redshift $z$ over the entire cosmological evolution, i.e., $z \\in [-1,\\infty)$. By using a theoretical algorithm of constructing the quintessence potential directly from the effective EoS parameter, we derive and discuss the general features of the resulting potential for this $w(z)$ function. Since the parameterization here discussed allows us to divide the parametric plane in defined regions associated to distinct classes of dark energy models, we use the most recent observations from type Ia supernovae, baryon acoustic oscillation peak and Cosmic Microwave Background shift parameter to check which class is observationally prefered. We show that the largest portion of the confidence contours lies into the region corresponding to a possible crossing of the so-called phanton divide line at some point of the cosmic evolution. ",
        "introduction": "Cosmological models with cold dark matter plus a $\\Lambda$ term ($\\Lambda$CDM) may explain most of the current astronomical observations (see, e.g., \\cite{review} for some recent reviews). However, from the theoretical viewpoint it is really difficult to reconcile the small value required by observations ($\\simeq 10^{-10} \\rm{erg/cm^3}$) with estimates from quantum field theories ranging from 50-120 orders of magnitude larger, as well as to explain why this is exactly the right value that is just beginning to dominate the energy density of the Universe today.  These issues make a complete cancellation of $\\Lambda$ (from some unknown symmetry of Nature) seem a plausible possibility and have also motivated a number of alternative explanations for the cosmic acceleration (for some of these alternative scenarios, see~\\cite{other, mad}). One of these possibilities, possibly the next simplest approach toward an accelerating model for the Universe, is to work with the idea that the dark energy component is due to a minimally coupled scalar field $\\phi$ which has not yet reached its ground state and whose current dynamics is basically determined by its potential energy $V(\\phi)$~\\cite{quint}. Clearly, however, such a procedure cannot provide a model-independent parameter space to be compared with the observational data.  Another way, widely explored in the literature,  is to build a phenomenological functional form for the dark energy equation of state (EoS), i.e., the ratio of its pressure to its energy density, $w \\equiv p/\\rho$, in terms of its current value $w_0$ and of its time-dependence $w_a =dw/d\\ln a$, and study its cosmological consequences as well as possible constraints on its behavior from observations. Usually, these parameterizations have not only the standard $\\Lambda$CDM  scenario ($w_0 = -1; w_a= 0$) but also the so-called $w$CDM model ($w_a = 0)$ as particular cases, so that constraints on its parameters may provide more accurate consistency checks to the original models.  Examples of  some EoS parameterizations recently explored are (see also~\\cite{para,para1,generalized}): \\begin{eqnarray} \\label{OtherParameterizations} w(z) = \\; \\left\\{ \\begin{tabular}{l} $w_0 + w_a z$ % \\, % \\quad \\quad  \\quad  \\quad \\quad \\quad \\quad \\quad  \\hspace{0.35cm}\\cite{linear}\\\\ \\\\ $w_0+w_az/(1+z)$ % \\quad \\quad  \\quad \\quad \\quad \\quad \\cite{cpl} \\\\ \\\\ $w_0-w_a\\ln(1+z)$ % \\quad \\quad  \\quad \\quad \\quad \\quad \\cite{efs} \\end{tabular} \\right. \\end{eqnarray} An interesting aspect worth mentioning is that it is difficult to obtain the above parameterizations from usual scalar field dynamics since they are not limited functions, i.e., the EoS parameter does not lie in the interval $w \\in[-1,1]$. In other words, this amounts to saying that when extended to the entire history of the universe, $z\\in [-1,\\infty)$, the three parameterizations above are divergent functions of the redshift (see also \\cite{para1} for a discussion)\\footnote{The $w(z)$ parameterizations given by Eq. (1) can be expressed as a single and generalized EoS function, i.e.,  $w(z) = w_0 - w_a\\frac{(1+z)^{-\\beta} -1}{\\beta}$, where the parameter $\\beta$ takes, respectively, the values $\\pm 1$ and $0$~\\cite{generalized}.}.  In Ref.~\\cite{edesio} we investigated some cosmological consequences of a new phenomenological parameterization: \\begin{equation} \\label{MyParameterization} w(z)=w_0+w_a\\frac{z(1+z)}{1+z^2}\\;, \\end{equation} or, equivalently,% \\begin{equation} \\label{MyParameterization1} w(z)=w_0+w_a\\frac{1-a}{1 - 2a + 2a^2}\\;, \\end{equation} which has the same linear behavior at low redshifts presented by the parameterizations discussed above but with the advantage of being a limited function of $z$ throughout the entire history of the Universe (see also~\\cite{pavon} for a recent analysis of a coupled quintessence model driven by a dark energy component parameterized by (\\ref{MyParameterization})).  Our goal in this paper is twofold. First, to derive a scalar field description for the dark component whose EoS parameter is given by Eq. (\\ref{MyParameterization}). To that end, we use the theoretical method of constructing the dark energy potential $V(\\phi)$  directly from the effective EoS, as developed in Ref.~\\cite{method}. We also generalize this algorithm to include the so-called phantom case. Second, to place constraints on the ($w_0$, $w_a$) plane of parameterization (\\ref{MyParameterization}) from current observational data to check which class of dark energy is observationally prefered. We use some of the most recent type Ia supernovae (SNe Ia) observations, namely, the so-called Union2 sample of 557 events~\\cite{union}, the  nearby + SDSS + ESSENCE + SNLS + Hubble Space Telescope (HST) set of 288 SNe Ia discussed in Ref.~\\cite{sdss} (which we refer to as SDSS compilation). We consider two sub-samples of this latter compilation that use SALT2~\\cite{salt2} and MLCS2k2~\\cite{mlcs2k2} SN Ia light-curve fitting method. Along with the SNe Ia data, and  to help break the degeneracy between the dark energy parameters we also use  measurements of the baryonic acoustic oscillation (BAO) peak at $z = 0.2, 0.35$ and 0.6~\\cite{bao,percival,blake} and the current estimate of the CMB shift parameter of Ref.~\\cite{wmap}.   \\begin{figure*}[t] \\centerline{\\psfig{figure=qf.ps,width=3.4truein,height=3.0truein,angle=270} \\hspace{0.5cm} \\psfig{figure=pf.ps,width=3.4truein,height=3.0truein,angle=270} \\hskip 0.1in} \\caption{Scalar field description of parameterization (\\ref{MyParameterization}) for three selected points in the quintessence (Panels a and b) and phantom (Panels c and d) regions.} \\end{figure*} ",
        "conclusions": "We have examined theoretical and observational aspects of the EoS parameterization given by Eq. (\\ref{MyParameterization})~\\cite{edesio}. By following the method of constructing the quintessence potential directly from the effective equation of state function developed in Ref.~\\cite{method}, we have derived the scalar field description for this $w(z)$ parameterization and extended our results for the case of phanton fields in which $w(z) < -1$. We also have performed a joint statistical analysis involving some of the most recent cosmological measurements of SNe Ia, BAO peak and CMB shift parameter. From a pure observational perspective, we have shown that both quintessence and phantom behaviours are acceptable regimes. In agreement with recent analyses, it has been shown that the largest portion of the confidence contours arising from these observations lies in the region of models that have crossed or will eventually cross the so-called phanton divide line at some point of the cosmic evolution."
    },
    "1107/1107.5714_arXiv.txt": {
        "abstract": "In a previous study published in {\\it Astrobiology}, we focused on the evolution of habitability of a 10~$M_\\oplus$ super-Earth planet orbiting a star akin to the Sun.  This study was based on a concept of planetary habitability in accordance to the integrated system approach that describes the photosynthetic biomass production taking into account a variety of climatological, biogeochemical, and geodynamical processes.  In the present study, we pursue a significant augmentation of our previous work by considering stars with zero-age main sequence masses between 0.5 and 2.0~$M_\\odot$ with special emphasis on models of 0.8, 0.9, 1.2 and 1.5~$M_\\odot$. Our models of habitability consider again geodynamical processes during the main-sequence stage of these stars as well as during their red giant branch evolution.  Pertaining to the different types of stars, we identify so-called photosynthesis-sustaining habitable zones (pHZ) determined by the limits of biological productivity on the planetary surface.  We obtain various sets of solutions consistent with the principal possibility of life.  Considering that stars of relatively high masses depart from the main-sequence much earlier than low-mass stars, it is found that the biospheric life-span of super-Earth planets of stars with masses above approximately 1.5~$M_{\\odot}$ is always limited by the increase in stellar luminosity.  However, for stars with masses below 0.9~$M_{\\odot}$, the life-span of super-Earths is solely determined by the geodynamic time-scale.  For central star masses between 0.9 and 1.5~$M_{\\odot}$, the possibility of life in the framework of our models depends on the relative continental area of the super-Earth planet. ",
        "introduction": "A central topic of contemporaneous astrobiology is the study of habitability around different types of stars, particularly main-sequence stars. This is motivated by the fact that stars spent most of their lifetimes on the main-sequence \\citep[e.g.,][]{maed88}.  Furthermore, it is noteworthy that the distribution of stellar spectral types is strongly tilted toward low-mass stars, i.e., K and M-type stars, as implied by detailed determinations of the initial mass function in the solar neighborhood and beyond \\citep[e.g.,][]{krou02,chab03}. For the present study we want to consider stars of large abundance in the galactic disk, within a reasonably large mass range, and which all have the potential of being central objects to habitable planetary systems.  Thus, it is the aim of the present paper to significantly augment our previous study that dealt with the evolution of circumstellar habitability for stars akin to the Sun ($1~M_\\odot$), see \\cite{bloh09}, by considering stars with masses between 0.5 and 2.0~$M_\\odot$.  Moreover, we will adopt a wider grid of planetary masses, ecompassing 1 to 10 Earth masses ($M_\\oplus$).  \\cite{bloh09} in Paper~I presented a detailed thermal evolution model for a 10 Earth-mass planet in the environment of a star like the Sun while also adopting an up-to-date solar evolutionary model \\citep{schr08}.  The planetary model of habitability has been based on the integrated system approach, which describes the photosynthetic biomass production taking into account a variety of climatological, biogeochemical, and geodynamical processes.  This allowed us to identify a so-called photosynthesis-sustaining habitable zone (pHZ) determined by the limits of biological productivity on the planetary surface.  The model also considered the principle possibility of habitability during stellar evolution along the Red Giant Branch (RGB).  It was found that the solar pHZ increases in width over time and moves outward, as expected.  For example, for ages of 11.0, 11.5, 12.0, and 12.1 Gyr, the pHZ is found to extend from 1.41 to 2.60, 1.58 to 2.60, 4.03 to 6.03, and 6.35 to 9.35 AU, respectively.  The approach of evaluating habitability based on assessing the pHZ is alternative to the study of the climatological habitable zone used by \\cite{kast93} and others. Note that concerning Earth-mass planets, a detailed study of geodynamic habitability based on the pHZ was previously presented by \\cite{fran00b}. They found that Earth will be rendered uninhabitable after 6.5~Gyr as a result of plate tectonics, notably the growth of the continental area (enhanced loss of atmospheric CO$_2$ by the increased weathering surface) and the dwindling spreading rate (diminishing CO$_2$ output from the solid Earth).  This work already considered systems of different types of main-sequence stars; however, it was based on an earlier version of stellar evolution models given by \\cite{scha92}.  Moreover, it also did not include stellar post--main-sequence evolution.  The study by \\cite{fran00b} implies that there is no merit in investigating the future habitability of Earth during long-term stellar evolution, as in the framework of pHZ models, the lifetime of habitability is limited by terrestrial geodynamic processes.  However, this situation is expected to be significantly different for super-Earth planets due to inherent differences compared to Earth-mass planets \\citep[e.g.,][]{vale07}. A further motivation for this type of work stems from the ongoing discovery of super-Earths in the solar neighborhood with the Gliese 876 \\citep{rive05} and Gliese 581 \\citep{udry07,vogt10} systems as prime examples. In the following, we describe aspects of stellar evolutionary processes pertaining to our study.  Next, we discuss the definition of the photosynthesis-sustaining habitable zone, including the relevant geodynamic assumptions.  Thereafter, we present our results and discussion.  Finally, we convey our summary and conclusions. ",
        "conclusions": "We investigated the habitability of super-Earth planets in the environments of main-sequence stars of masses between 0.5 and and 2.0~$M_\\odot$.  This work complements a previous study that has been focused on star akin to the Sun ($M = 1 M_\\odot$), see \\cite{bloh09}.  Both studies employ the concept of planetary habitability based on the integrated system approach that describes the photosynthetic biomass production taking into account a variety of climatological, biogeochemical, and geodynamical processes. Among the various assumptions entertained as part of our study, we consider that super-Earth planets are expected to develop plate tectonics, an assumption critical to our models, although it is currently under scrutiny.  Pertaining to the different types of stars, we identify so-called photosynthesis-sustaining habitable zones (pHZ) determined by the limits of biological productivity on the planetary surface.  We obtain various sets of solutions consistent with the principal possibility of life.  Our models are based on advanced version of the Eggleton stellar evolution code that considers updated opacities and an updated equation of state as described by \\cite{pols95,pols98}. Additionally, an improved description of mass loss is implemented following the work by \\cite{schr05}.  Considering that high-mass stars depart from the main-sequence much faster than low-mass stars, it was found that the biospheric life-span of super-Earth planets of stars with masses above approximately 1.5~$M_{\\odot}$ is always limited by the increase in stellar luminosity dictated by stellar nuclear evolution.  However, for stars with masses below 0.9~$M_{\\odot}$, the life-span of super-Earths it is solely determined by the geodynamic time-scale.  Clearly, the luminosity of stars with spectral type mid-K or later has not deviated much from their ZAMS luminosities owing to the limited age of the Universe.  For central star masses between 0.9 and 1.5~$M_{\\odot}$, the possibility of life in the framework of our models depends on the relative continental area of the super-Earth planets.  A crucial finding of our study is that planets of large geological age are expected to be water worlds, i.e., planets with a relatively small continental area.  The reason why land worlds are considered to be unfavorable is that planets with a considerable continental area have relatively high weathering rates that provide the main sink of atmospheric CO$_2$.  Therefore, this type of planets are unable to build up CO$_2$-rich atmospheres that prevent the planet from freezing or allowing for the possibility of photosynthesis-based life.  Our study is also both expanding and superseding previous work, including the study by \\cite{fran00b}.  This work solely considered Earth-mass planets instead of also including super-Earth planets.  Additionally, it considered a linear growth model for the continental growth and has refrained from taking into account post-main sequence evolution.  Although the evolutionary time-scales for stars on the RGB are short compared to the time spent on the main-sequence, habitable planets around red giants should be considered a realistic possibility, as previously pointed out by \\cite{lope05}, \\cite{bloh09}, and others.  However, for central stars with higher masses than the Sun, a more rapid evolution will occur that will also place the significant temporal and spatial constraints on planetary habitability when the central stars have reached the RGB.  \\pagebreak"
    },
    "1107/1107.5522_arXiv.txt": {
        "abstract": "{ We investigate in detail the cosmological consequences of realistic angular dependent potentials in the brane inflation scenario. Embedding a warped throat into a compact Calabi-Yau space with all moduli stabilized breaks the no-scale structure and induces angular dependence in the potential of the probe D3-brane. We solve the equations of motion from the DBI action in the warped deformed conifold including linearized perturbations around the imaginary self-dual solution. Our numerical solutions show that angular dependence is a next to leading order correction to the dominant radial motion of the brane, however, just as angular motion typically increases the amount of inflation (spinflation), having additional angular dependence also increases the amount of inflation. We also derive an analytic approximation for the number of e-foldings along the DBI trajectory in terms of the compactification parameters. } ",
        "introduction": "It is now generally accepted that the very early Universe underwent a period of rapid expansion, known as inflation, resulting in a very nearly flat, homogeneous and isotropic initial state. While a simple scalar field model of inflation with a suitable potential satisfies many of the cosmological requirements, such models are rather ad hoc from the high energy particle theory point of view. The challenge is to find a theory which has a clear derivation from a fundamental high energy theory incorporating gravity at the quantum level. String theory naturally is the prime candidate for such a fundamental theory, but one of the major challenges has been to show that in a non-vanishing fraction of the vast number of string vacua the inclusion of various compactification effects in the effective field theory leaves a suitable inflationary model.  In recent years there has been significant progress within the `KKLT' framework, \\cite{KKLT}, in which the older idea of brane inflation \\cite{brinfl}, is realised via the motion of a brane in internal, hidden, extra dimensions \\cite{KKLMMT} (see also \\cite{followon,BDKKM,tip,tipP,IIB} and the reviews \\cite{reviews}). In these scenarios the inflaton potential (of the position of a probe D3-brane) is flattened by fine-tuned cancellation of correction terms from moduli stabilizing wrapped branes and bulk effects so that inflation can occur. However, whether or not the potential can be made flat to meet the slow-roll conditions, DBI inflation, \\cite{ST}, is also possible, even when the potential is steep. In this case, while the motion of the brane can be strongly relativistic (in the sense of a large $\\gamma$-factor) strong warping of the local throat region renders their contribution to the local energy density subdominant to that of the inflaton potential terms.  In either case, the approach taken is to consider a D3-brane moving in a warped throat region of a Calabi-Yau flux compactification of type IIB theory with ISD conditions \\cite{GKP}. However, when the UV end of the warped throat is attached to the compact Calabi-Yau space with all moduli stabilized, violations of ISD conditions with important implications for the action of the brane are expected. In particular, the potential of a mobile D3-brane in the compactified throat geometry receives angular dependent corrections, \\cite{BDKKM}, which until recently, have been largely neglected (although see \\cite{tip,tipP,ABMX}). In slow roll inflation, it was presumed that the angular directions would stabilise rapidly, with the radial (slow-roll) direction dominating the inflationary trajectory. However, for generic brane motion the effect of angular motion is less clear, particularly in the presence of angular terms in the potential.  Angular motion of branes was initially explored in the probe limit, i.e.\\ where the brane does not back-react at all on either the internal or external dimensions. Unsurprisingly, the angular motion has a conserved momentum, which can give interesting brane universes (see e.g.\\ \\cite{braneon,EGTZ, sling}), however the mirage style, \\cite{mirage}, interpretation of the cosmology of these universes leads to a rather unsatisfactory picture. Based on the probe understanding, it was conjectured that angular motion would not affect a more realistic inflationary scenario to any great extent, an expectation largely borne out by the ``spinflation'' study, \\cite{spin}, which found a marginal increase of a couple of e-foldings due to angular motion, coming mainly from the initial stages of inflation before the angular momentum becomes redshifted away. This increase is however parameter sensitive, a point not noted in this original study. In \\cite{spin}, a general DBI-inflationary universe was considered near the tip of the warped deformed conifold throat, the Klebanov-Strassler (KS) solution \\cite{KS}, with a simple radial brane potential; with a more realistic potential including angular terms, the spinflationary picture could potentially be rather different.  In this paper therefore, we investigate the cosmological implications of including angular dependence in the DBI brane inflation scenario. Building on the results of Baumann \\emph{et al.} \\cite{BDKKM}, we consider D3-brane motion in the warped throat region of the compact Calabi-Yau subject to UV deformations of the geometry that induce angular dependent corrections in the potential of the probe D3-brane. Taking into account linearized perturbations around the ISD solution, we solve the D3-brane equations of motion from the DBI action with angular dependence induced by the leading correction to the potential allowed by the symmetries of the compactification. Our aim is to consider angular momentum in a fully UV/IR consistent fashion, and to account for angular momentum in a more general potential. As with the simpler radial spinflation potential, our numerical solutions show that angular dependence tends to increase the inflationary capacity of a trajectory, increasing the number of  e-foldings, albeit at a subdominant level. To a large extent, the trajectories are still predominantly radial, however, do exhibit rotational motion due to the angular potential.  An important question is how generic such trajectories are. In general, as the brane arrives in the throat, one expects a range of initial conditions in terms of angular values and velocities. We find that the D brane trajectories and number of e-foldings are dependent more on model parameters than on the initial condition of the brane motion, thus indicating that the results of our investigation are reasonably robust. ",
        "conclusions": "In this paper we studied generic brane motion in a warped deformed throat incorporating full angular dependence of both the brane potential and brane motion. This is the first such study of D3-brane inflation which samples the many features of a warped throat, including both UV and IR features of the geometry, as well as IR consistent supergravity corrections. We considered a generic mass term for the D3 brane potential as well as an analytic linearized perturbation around the ISD background, which to our knowledge is the first closed form analytic radial eigenfunction on the KS background.  Our results show that angular features of the brane motion are dependent on the compactification parameters more than on initial conditions: For strongly deformed throats (relatively large $\\epsilon$) the brane `sees' the full rounded structure of the tip and has a very rich angular motion. These oscillatory trajectories however have very little cosmological expansion. Conversely, for very sharp throats ($\\epsilon\\ll1$) the brane enters a strongly DBI inflating motion, which is geometry dominated and mainly radial in nature. The coordinate velocities of the brane are very small, and the brane falls to the minimum of the potential in its first sweep down the throat, only oscillating right at the exit of inflation and at the minimum of the potential. The models and initial conditions considered were deliberately multi dimensional, and the brane samples the full IR region of the throat.  In this analysis, we have deliberately considered a set-up in which there is no obvious slow-roll regime. In particular, using the correct, fully infrared consistent corrections to the potential, we see that an inflection point potential is no longer consistent. Our model explicitly follows brane inflation from the UV to the IR region of the throat, and as such is not particularly suitable for angular inflation at the tip, as the potential deep in the IR is not sufficiently steep. However, the corrections we considered have also been used in a restricted context in an angular tip inflation, \\cite{tipP}, although in this case, the full dependence of the eigenfunction on the radial direction was not considered.  In all cases, these DBI trajectories have large $\\gamma$ factors and hence will generically generate large non-gaussianities if used as pure inflation models in their own right, \\cite{ST,NG}. In order to construct a viable inflationary model therefore, some alternative mechanism must be found to produce perturbations, or this motion must be part of a bigger inflationary picture in which CMB scale fluctuations have already been produced (see e.g.\\ \\cite{tipP}). This paper has focussed on seeking a UV/IR consistent inflationary picture, thus we have not explored the full detail of cosmological perturbation theory, however, there are several options which could, in principle be incorporated into this model.  One particularly interesting possibility is that some curvaton mechanism might arise, either from one of the additional angular degrees of freedom, \\cite{DBIc} or from the vector excitations inherent in this type of model, \\cite{DWZ}. In the curvaton picture, \\cite{curv}, a (possibly more natural) option is presented in which additional fields generate perturbations, rather than having the inflaton performing all roles.  Because our angular potential is explicit, it clearly does not depend on several of the other internal angles, which could therefore provide flat directions in the potential. However, since these additional scalar degrees of freedom are part of the multifield sigma model of the throat, the dynamics of perturbation theory is quite involved and only a full analysis would reveal if this is indeed possible. Perhaps a more promising and natural approach which has recently been explored in \\cite{DWZ} is that perturbations might be generated via a vector curvaton, although in that case the authors did not explore angular motion or corrections to the supergravity potential. This would be an extremely interesting avenue to explore.  Another interesting observation is that a nonminimal gravitational coupling (depending on the non-radial degrees of freedom) tends to increase the inflationary capacity, while decreasing the $\\gamma$-factor, \\cite{vdb}. However, the analysis in \\cite{vdb} was from a more empirical point of view, more reminiscent of a k-inflation model, \\cite{Kinf}, and it would be interesting to see whether similar results are attainable in an explicit supergravity model arising from higher order corrections to the solution (e.g.\\ coupling to the Ricci scalar) considered in this paper.  It is also possible that since the $\\gamma$ factor varies throughout the motion, for some carefully chosen parameters this problem could be circumvented. It is interesting to use the intuition gleaned here to consider the late time evolution of more conventional slow-roll brane models, and explore what might happen as the brane approaches the tip. A slow-roll model such as the delicate universe, \\cite{BDKKM}, uses supergravity perturbations which are strictly only allowed in the pure $T^{1,1}$ throat, however, setting this aside for a moment, we can speculate on the final brane downfall to the tip. A vast scanning of parameter space for potentials is not necessary, as the angular motion and dependence is more a quantitative than qualitative factor, and our investigation has given insight into the effect of the different compactification parameters as they vary.  Typically, a slow roll model of brane inflation will be in the intermediate adS regime of the throat geometry, and thus requires a small deformation parameter. However, as the deformation parameter becomes very small, the final roll to the tip can take a long time, and be highly relativistic. The compactification data are generally expressed differently in references \\cite{BDKKM,ABMX}, in terms of D3 flux, $N$, the warp hierarchy between the inflating and IR region, $a_0 = e^{A(0)}$, however, knowledge of $M_p$, $T_3$ and the UV cut-off scale of the throat, $r_{UV}$, allows a translation between the different parametrizations. For example, to compare to the study of Agarwal et al., \\cite{ABMX}, saturating the $M_p$ bound and taking their stated parameter values leads to deformation parameters of $\\epsilon \\sim 10^{-4}-10^{-7}$ for $a_0\\sim 10^{-3}-10^{-5}$. With such small deformation parameters, the final sweep to the throat tip could easily incorporate several to many e-foldings of DBI inflation, and raises the question of just how much impact this final sweep could have.  Finally, our analysis has focussed on motion in only one angular direction. This was primarily to allow clarity of the angular terms in the potential, but also for simplicity. We expect that including fully generic angular motion (which would significantly complicate the analysis) will have an effect subdominant to the subdominant angular motion presented here, however, it would be useful to confirm this, in particular to confirm or rule out the possibility of a curvaton arising in this sector. In addition, our analysis could be further extended by taking higher order corrections into account, which will also have angular dependence. While we would expect that this would be sub-dominant to the linearized correction, it would be interesting to see precisely what the impact is of these detailed corrections to the inflaton action."
    },
    "1107/1107.2244_arXiv.txt": {
        "abstract": "{In order to study the association of the origin of ultra high energy cosmic rays (UHECR) with active galactic nuclei (AGN) at all levels of their activity we require an unbiased sample of black holes.}{Here we describe such a sample, of about 6 000 black holes, within the local Universe, inside the GZK (Greisen Zatsepin Kuzmin) limit, around 100 Mpc.}{ The starting point is the 2 micron all sky survey, with the next steps as: test its completeness down to low flux densities, confine it to redshifts $z < 0.025$, limit it to early Hubble type galaxies, test with B-V colors and with the FIR/Radio ratio the possible separation in classes of sources for the ultra high energy cosmic rays, use the spheroidal stellar component - black hole mass relationship to derive black hole masses, and test them with known black hole masses.}{The statistics are consistent with the mass function of black hole masses, with a relatively flat distribution to about $10^{8} \\, M_{\\odot}$, and thereafter a very steep spectrum. Our sample cuts off just below $10^{7} \\, M_{\\odot}$, indicating a gap for lower masses, below somewhere near $10^{6} \\, M_{\\odot}$. We construct the sky distribution for all black holes above $10^{7} \\, M_{\\odot}$, $10^{8} \\, M_{\\odot}$, and $3\\times 10^{8} \\, M_{\\odot}$, and also show the sky distribution for the five redshift ranges 0.0 - 0.005, 0.005 - 0.0010, 0.010 - 0.015, 0.015 - 0.020, and 0.020 to 0.025.}{} ",
        "introduction": "The samples of active galactic nuclei are small in the nearby universe and incompletely distributed on the sky, for example, there are only 865 objects in the V{\\'e}ron catalog inside a 100 Mpc region around our Galaxy \\citep{2010A&A...518A..10V}. As many models developed for the determination of the origin of UHECR require information concerning the past activity of a black hole, and little of the current activity, (see, e.g.  \\citet{1999MNRAS.307..491B,2000MNRAS.316L..29B,2002APS..APRB17041B,2002APh....16..265L,2002PhRvD..66g4014M,2008MNRAS.383..663B,1977MNRAS.179..433B,2007APh....27..473L}) it is useful to derive a nearly complete, all sky sample of black holes in the nearby Universe, inside the sphere around our Galaxy determined by the GZK distance, \\citet{1966PhRvL..16..748G,1966ZhPmR...4..114Z}.  We note that in fact all black holes can be detected as active, some at very low level, most easily via non-thermal radio emission \\citep{1984A&A...130L..13P,1984ApJ...283..479E,2000ApJ...542..186N}. One very good example is the black hole at the center of our Galaxy, it\u2019s past activity was discovered by looking at the entire electromagnetic emission \\citep{2010PNAS..107.7196M}.  The overall strategy of finding the associations between the astrophysical sources and the events observed in the sky is to develop a nearly complete sample of possible sources, separate them in sub-samples according to their nature (e.g. blue or red star dominated, starburst or non-starburst galaxies), take the maximal sources from each set and model the physical processes that are involved in the production of the energetic particles in order to obtain the maximal flux of the cosmic rays. Then the next step is to model and implement the propagation of these particles from the distribution of sources to their arrival distribution here on Earth including the deflection in a structure of magnetic fields. Finally one needs to make the statistical correlation between the virtual events and the real observed ones. In this article the first steps are taken with the next ones treated in the following articles. ",
        "conclusions": "The sample constructed here will allow a proper statistical treatment of the possible association between ultra high energy cosmic rays and the activity of black holes being an all sky complete collection of galaxies.  This kind of correlations are also useful in testing physical theoretical models that are responsible for the birth and evolution of massive and super-massive black holes.  Also this catalog of massive black hole can be use in statistical interpretation of the distribution of black holes with various applications from indicating the nature of dark matter to tracing the activity of nearby classes of galaxies.   \\acknowledgement{Work with PLB was supported by contract AUGER 05 CU 5PD 1/2 via DESY/BMB and by VIHKOS via FZ Karlsruhe; by Erasmus/Sokrates EU-contracts with the universities in Bucharest, Cluj-Napoca, Budapest, Szeged, Cracow, and Ljubljana; by the DFG, the DAAD and the Humboldt Foundation; also by research foundations in Korea, China, Australia, India and Brazil. L.I.C. was supported by a research program from the LAPLAS 3 and CNCSIS Contract 539/2009. LIC wishes to express his thanks to the MPIfR for support when finishing this project, also discussions with A. Caramete are acknowledged. 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. This research also made use of the ViZier system at the Centre de Donne{\\'e}s astronomiques de Strasbourg (CDS)  \\citep{2000A&AS..143...23O}. We also acknowledge the usage of the HyperLeda database (http://leda.univ-lyon1.fr).}"
    },
    "1107/1107.2591_arXiv.txt": {
        "abstract": "We present our results on  the structure and activity of massive galaxies at $z=1-3$ using one of the largest (166 with $M_\\star\\geq5\\times10^{10}$ $M_\\odot$) and most diverse samples of massive galaxies derived from the GOODS-NICMOS survey: (1)~S\\'ersic fits to deep NIC3/F160W images indicate that the rest-frame optical structures of  massive galaxies are very different at $z=2-3$ compared to $z\\sim0$. Approximately 40\\% of massive galaxies are ultra-compact ($r_e\\le2$ kpc), compared to less than 1\\% at $z\\sim0$. Furthermore, most ($\\sim65$\\%)  systems at $z=2-3$ have a low S\\'ersic index $n\\le2$,  compared to $\\sim13\\%$ at $z\\sim0$. We present evidence that the $n\\le2$ systems  at $z=2-3$  likely contain prominent disks, unlike most  massive $z\\sim0$ systems. (2)~There is a correlation between structure and star formation rates (SFR). The majority ($\\sim85$\\%)  of non-AGN massive galaxies at $z=2-3$,  with SFR high enough to yield a 5$\\sigma$ (30$\\mu$Jy) 24 $\\mu$m $Spitzer$ detection have low $n\\leq 2$. Such $n\\leq 2$ systems host the highest SFR. (3)~The frequency of AGN is $\\sim40\\%$ at $z=2-3$. Most  ($\\sim65\\%$) AGN hosts have disky ($n\\leq2$) morphologies. Ultra-compact galaxies  appear quiescent in terms of both AGN activity and star formation. (4)~Large stellar surface densities imply massive galaxies at $z=2-3$ formed via rapid, highly dissipative events at $z>2$. The large fraction of $n\\le2$ disky systems suggests cold mode accretion complements gas-rich major mergers at $z>2$. In order for massive galaxies at $z=2-3$ to evolve into present-day massive E/S0s, they need to significantly increase ($n$, $r_e$). Dry minor and major mergers may play an important role in this process. ",
        "introduction": "Studies of high-redshift galaxies are essential for testing and constraining models of galaxy formation.  Conventional wisdom suggests galaxies are assembled and shaped by a combination of mergers, smooth accretion, and internal secular evolution. Galaxies form inside cold dark matter halos that grow hierarchically through mergers with other halos and gas accretion (Somerville \\& Primack(1999); Cole et al. 2000; Steinmetz \\& Navarro 2002; Birnboim \\& Dekel 2003; Kere{\\v s} et al. 2005; Dekel \\& Birnboim 2006; Dekel et al. 2009a; Dekel et al. 2009b; Kere{\\v s} et al. 2005;  Kere{\\v s} et al. 2009; Brooks et al. 2009; Ceverino et al. 2010), while  internal secular evolution (Kormendy \\& Kennicutt, 2004; Jogee et al. 2005)  redistributes accreted material. Within the paradigm of hierarchical assembly, a number of issues remain. It is not known when and how the main baryonic components of modern galaxies (bulges, disks, and bars) formed, but  the global stellar mass density rose substantially between $z\\sim1-3$, reaching $\\sim50\\%$ of its present value by $z\\sim1$ (Dickinson et al. 2003b; Drory et al. 2005; Conselice et al. 2007; Elsner et al. 2008; P{\\'e}rez-Gonz{\\'a}lez et al. 2008).  It is also not clear how high-redshift galaxies evolve into present-day galaxies.  Complex baryonic physics such as mergers, gas dissipation, and feedback are all at work to an extent. There is also mounting evidence that cold-mode accretion (Birnboim \\& Dekel 2003; Kere{\\v s} et al. 2005; Dekel \\& Birnboim 2006; Dekel et al. 2009a; Dekel et al. 2009b; Kere{\\v s} et al. 2005; Kere{\\v s} et al. 2009; Brooks et al. 2009; Ceverino et al. 2010) is important for building  star-forming galaxies. This process is particularly effective in galaxies with halos of mass below $10^{12}$ $M_\\odot$ such that cold-mode accretion dominates the global growth of galaxies at high redshifts and the growth of lower mass objects at late times.  High-redshift galaxies are different from local galaxies. Within the framework of hierarchical assembly, early, high-redshift galaxies are expected to be smaller, at a given mass, than their present-day counterparts. The size difference is predicted to be a factor of a few at $z=2-3$ (Loeb \\& Peebles 2003; Robertson et al. 2006; Khochfar \\& Silk 2006; Naab et al. 2007).  Several recent studies using rest-frame optical data provide evidence for size evolution among massive galaxies (Guzman et al. 1997; Daddi et al. 2005; Trujillo et al. 2006, 2007; Zirm et al. 2007; Toft et al. 2007; Longhetti et al. 2007; Cimatti et al. 2008; Buitrago et al. 2008; van Dokkum et al. 2008, 2010; van der Wel et al. 2011). Aside from size evolution, there is some evidence that the nature of red galaxies changes at higher redshift. At $z\\lesssim 1$, the red sequence primarily consists of old, passively evolving galaxies (Bell et al. 2004). Among extremely red galaxies (EROs) at $z=1-2$, less than 40\\% are morphologically early types (Yan \\& Thompson 2003; Moustakas et al. 2004). It is well known that star formation rates were more intense at higher redshift (Daddi et al. 2007; Drory \\& Alvarez 2008), and a link has been found between star formation, size, and morphology at $z\\sim2.5$. Toft et al. (2007) and Zirm et al. (2007) find from NICMOS rest-frame optical imaging that blue star-forming galaxies are significantly more extended than red quiescent galaxies. Additionally, examples of rapidly star-forming galaxies (SFR $\\sim50-200$ $M_\\odot$ yr$^{-1}$) at $z\\sim2-3$, whose ionized gas kinematics are consistent with turbulent rotating disks, are found in the SINS survey (F{\\\"o}rster Schreiber et al. 2009; Genzel et al. 2008; Shapiro et al. 2008).  Progress on understanding the evolution of massive galaxies at high redshift has been hindered by significant observational challenges. The deep optical surveys carried out by $HST$ $ACS$, such as the Hubble Ultra Deep Field (HUDF, Beckwith et al. 2006) and the Great Observatories Origins Deep Survey (GOODS, Giavalisco et al. 2004), trace rest-frame optical galaxy morphology only out to $z\\sim1$. At higher $z$, bandpass shifting effects cause filters to trace progressively bluer bands, and optical filters trace rest-frame UV at $z\\gtrsim2$. UV light traces massive young stars, but manages to set few constraints about the overall mass distribution, making it difficult to probe the structure and mass of galaxy components at early epochs.  Without high-resolution, deep, rest-frame optical imaging, it is not possible to robustly compare structural parameters in galaxies across redshift. NIR imaging is required to probe the rest-frame optical at $z\\sim1-3$. Unfortunately, deep NIR imaging with $HST$ has been completed for a limited number of galaxies over relatively small fields and small volumes at $z>1$, with most pointings being within the Hubble Deep Fields and the Hubble Ultra Deep Field due to the inefficiency of the NICMOS camera in covering large areas (e.g., Dickinson et al. 2004; Thompson et al. 2005; Zirm et al. 2007; van Dokkum et al. 2008). While ground-based NIR imaging surveys (e.g., Kajisawa et al. 2006; Retzlaff et al. 2010) efficiently cover wide fields at resolutions almost comparable to $HST$ NICMOS, the depths reached are at least an order of magnitude shallower.  A large area, high-resolution, deep, space-based NIR survey would be bountiful for galaxy formation studies.  The GOODS-NICMOS Survey (GNS; Conselice et al. 2011), covering 44 arcmin$^{2}$ of the GOODS fields with NIC3, is a strong first effort. The GOODS-North and GOODS-South are among the best-studied regions in the sky and are a natural choice for such a survey. The GOODS fields already have deep data from $HST$ $ACS$ (Giavalisco et al. 2004), $Spitzer$ IRAC/MIPS (Dickinson et al. 2003a), and Chandra (Giacconi et al. 2002; Alexander et al. 2003; Lehmer et al. 2005; Luo et al. 2008), among others. GNS consists of 60 pointings centered on massive ($M_\\star>10^{11}$ $M_\\odot$) galaxies at $z>2$, observed to a depth of $H=26.8$ magnitudes. The value of GNS lies in the fact that the target fields were optimized to include massive galaxies selected by multiple methods in order to create an unbiased sample (see Conselice et al. 2011). There are additional massive galaxies in each field beyond the 60 main targets, so that there are 82 galaxies with $M_\\star \\geq 10^{11}$ $M_\\odot$ at $z=1-3$ across all pointings. Thus, the GNS data contain one of the largest samples of very massive galaxies at high redshift with rest-frame optical imaging, and they robustly probe massive galaxies when the Universe was less than 1/3 of its current age, during the epoch of bulge and disk formation.  The goal of this work is to investigate the evolution of massive galaxies over $z=1-3$ with this unique sample. We take advantage of the existing rich ancillary data to derive star formation rates (SFR) from 24 $\\mu$m detections and look for AGN activity based on X-ray detections and mid-IR SEDs. We correlate rest-frame optical structural parameters with SFR to gain insight into how massive galaxies are expected to evolve.  The plan of this paper is as follows.  We discuss the data and sample properties in \\S\\ref{data}. In \\S\\ref{sprop} we describe the measurement of structural parameters, and in \\S\\ref{sressp} we make a detailed comparison with $z\\sim0$ galaxies of similar stellar mass. A detailed artificial redshifting experiment is conducted in \\S\\ref{redshift} to explore the impact of instrumental and redshift-dependent effects on structural parameters. In $\\S$\\ref{sfrate}, we measure star formation properties based on $Spitzer$ MIPS  24 $\\mu$m detections and discuss how they relate to structural properties. Estimates of  the mass and fraction of cold gas in massive star-forming galaxies  at $z=2-3$ are presented in   $\\S$\\ref{scold}. In $\\S$\\ref{agn}, we use a variety of techniques (X-ray properties, IR power-law, and IR-to-optical excess) to identify AGN and consider how galaxy activity relates to galaxy structure. Finally, in \\S\\ref{discuss} and \\S\\ref{summary}, we discuss and summarize our results. All calculations assume a flat $\\Lambda$CDM cosmology with $\\Omega_\\Lambda = 0.7$ and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We use  a variety of techniques (X-ray properties, IR power-law, IR-to-optical excess, and mid-IR colors) to identify Active Galactic Nuclei (AGN) among the massive GNS galaxies because selection based on X-ray emission alone may fail at high redshift in the case of Compton-thick AGN where much of the soft X-ray emission is Compton scattered or absorbed by thick columns of gas ($N_H \\gg 10^{24}$~cm$^{-2}$; Brandt et al. 2006). We briefly summarize here and in Table~\\ref{tabagnfrak} the number of AGN identified by each of the selection methods\\footnote{The mid-IR selection criteria of Lacy et al. (2004) and Stern et al. (2005) were investigated but considered unreliable. Contamination from high-redshift star-forming galaxies drastically reduces their accuracy (e.g., Donley et al. 2008).  Applying these methods at $z=1-3$ would add more false-positives than true AGN.}.  \\begin{enumerate} \\item X-ray counterparts to the massive GNS sources were searched for in the CDF-N and CDF-S catalogs of Alexander et al. (2003) and Luo et al. (2008), as well as the ECDF-S catalogs of Lehmer et al. (2005). A total of 33/166 massive GNS galaxies had counterparts within $1\\farcs5$ across all catalogs.  \\item Following Alonso-Herrero et al. (2006) and Donley et al. (2008), we look for AGN power-law emission over  $z=1-3$ using SEDs from the IRAC bands at 3.6, 4.5, 5.8, and 8.0 $\\mu$m.  The IRAC SEDs were fit with a power-law SED ($f_\\nu \\propto \\nu^{\\alpha}$). There are only 3/166 sources with power-law index $\\alpha \\le -0.5$ that are considered power-law galaxies (PLGs) and obscured AGN candidates.  \\item Heavily obscured AGN may be present in highly  reddened, IR-excess galaxies. Fiore et al. (2008) identify obscured AGN candidates in IR-bright, optically faint, red galaxies over $z=1.2-2.6$ using the criteria $f_{24\\mu m}/f_R \\geq 1000$ and $R-K>4.5$. We search for such IR-bright, optically faint systems with  $f_{24\\mu m}/f_R > 1000$ in our sample of massive galaxies. $R$-band flux is determined by linear interpolation between the $ACS$ $V$ and  $i$-band fluxes. We find 25 sources meeting this criteria, of which 16 are new AGN candidates not identified via the above two methods. \\end{enumerate}   Among the 166 massive GNS galaxies at $z=1-3$, the AGN fraction is 49/166 or $29.5 \\pm 3.5\\%$. When the results are broken down in terms of redshift, the AGN fraction rises with redshift, increasing from $17.9 \\pm 6.1\\%$ at $z=1-1.5$ to $40.3 \\pm 8.8 \\% $ at $z=2-3$. The percentage of AGN among all massive GNS galaxies is higher than at $z\\sim1$, where it is reported that less than 15\\% of the total 24~$\\mu$m emission at $z <1$ is in X-ray luminous AGN (e.g.,  Silva \\etal 2004; Bell \\etal 2005; Franceschini \\etal 2005; Brand \\etal 2006).  \\subsection{Relation Between AGN Activity and Structure}\\label{agnstru} We summarize the properties of the AGN host candidates and discuss their implications in terms host galaxy structure.  Figure \\ref{agnnr} shows the single S\\'ersic  $n$ versus $r_e$. Most ($80.6 \\pm 7.9\\%$) of the AGN  hosts at $z=2-3$  have $r_e > 2$ kpc and are not ultra-compact. AGN appear to be found preferentially in the more extended galaxies. Indeed, at $z=2-3$, the AGN fraction in ultra-compact galaxies is $\\sim2.7$ times lower than in extended galaxies ($20.0 \\pm 16.3\\%$ versus $53.2 \\pm 10.0\\%$). At $z=1-2$ the deficiency is a factor of $5.6$. Thus, the ultra-compact galaxies are more quiescent in terms of both AGN activity and SFR activity (see \\S\\ref{sfrate}).  Furthermore, a significant fraction of these AGN ($64.6 \\pm 10.7\\%$) have disky ($n\\leq2$) morphologies. Over half  ($58.2 \\pm 11.6\\%$) of the AGN candidates are both disky and not ultra-compact.   Similar statistics apply over $z=1-2$. The disky nature of AGN hosts at $1.5<z<3$ has been measured previously by Schawinski et al. (2011). From decomposition of the rest-frame optical light for 20 AGN imaged with $HST$ WFC3, they measure a mean S\\'ersic index of 2.54 and a mean effective radius of $3.16$ kpc.  Their results for ($n$, $r_e$) are consistent with our results for $z=2-3$ in Table~\\ref{tabagnfrak} and Figure \\ref{agnnr}. Furthermore, Kocevski et al. (2011, in prep.) find from visual classification of rest-frame optical morphologies that $51.4^{+5.8}_{-5.9}$ of X-ray selected AGN ($L_X \\sim 10^{42-44}$ erg s$^{-1}$) at $1.5<z<2.5$ reside in galaxies with visible disks; only $27.4^{+5.8}_{-4.6}$ reside in pure spheroids.  If the disky AGN host candidates  host massive black holes, then massive black holes are present in galaxies that are not dominated by a massive spheroid. In the local Universe,  nearly all massive galaxies are believed to host a central supermassive  black hole (Kormendy 1993; Magorrian et al. 1998;  Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Marconi \\& Hunt 2003), and the black hole mass is tightly  related to the bulge stellar velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt  et al. 2000). This has led to the suggestion that the black hole and bulge or spheroid probably grew in tandem (e.g., Cattaneo \\& Bernardi 2003; Hopkins et al. 2006). The presence at $z=2-3$ of luminous and potentially massive black holes in high mass galaxies that do not seem to have a prominent bulge or spheroid may be at odds with this picture.  \\label{summary} We present a study of the structure, activity, and evolution of massive galaxies at $z=1-3$ using  deep ($5\\sigma$ limiting magnitude of $H$=26.8 AB for an extended source of diameter $0\\farcs7$), high resolution   (PSF$\\sim 0\\farcs3$)  NIC3/F160W images from the GOODS-NICMOS Survey (GNS),  along with complementary ACS, $Spitzer$ IRAC and MIPS, and Chandra X-ray data. One of the strengths of our study is that the NIC3/F160W data provide rest-frame optical imaging over $z=1-3$ for one of the largest (166 galaxies with $M_\\star \\geq 5 \\times 10^{10}$ $M_\\odot$ and 82 with $M_\\star \\geq 10^{11}$ $M_\\odot$), most diverse, and relatively  unbiased samples of massive galaxies  at $z=1-3$  studied to date. Our main results are summarized below.  \\vspace{3mm} {\\it  1.  Structure of massive galaxies at rest-frame optical wavelengths: } We  analyze the rest-frame optical structure of the massive galaxies by fitting single S\\'ersic profiles to the 2D light distribution in the NIC3/F160W images. We find that the rest-frame optical structures of the massive galaxies are very different at $z=2-3$ compared to $z\\sim$~0, with  their  S\\'ersic index $n$ and half-light radius  $r_e$ being strikingly offset toward lower values compared to $z\\sim$~0. (Table~\\ref{tabmgc} and Figure~\\ref{fnre}). Through extensive tests and artificial redshifting experiments we conclude that the  offset in ($n$, $r_e$) between massive  galaxies at $z=2-3$ and   $z\\sim0$  is real and not primarily driven by systematic  effects related  to the fitting techniques instrumental effects, or  redshift-dependent effects (e.g., cosmological surface brightness dimming and the loss of spatial resolution).  In effect, we find {\\it a  large population of ultra-compact ($r_e \\le 2$~kpc) systems}, as well as {\\it a  dominant population of systems with low $n\\le2$ disky morphologies} at $z=2-3$. We further describe these populations below.  We find  that  approximately 40\\% ($39.0  \\pm 5.6$\\%  for  $M_\\star \\geq 5\\times 10^{10}$ $M_\\odot$ and $39.0  \\pm 7.6$\\%  for  $M_\\star \\geq 1\\times 10^{11}$ $M_\\odot$) of the massive  galaxies  at $z = 2-3$ are in the form of ultra-compact ($r_e \\le 2$~kpc)  galaxies compared to less than 1\\% at $z\\sim$ 0 (Table~\\ref{tabmgc} and Figure~\\ref{fnre}). These ultra-compact galaxies are practically unmatched among $z\\sim 0$ massive galaxies,  and their surface brightness in the rest-frame optical can be 4-6  magnitudes brighter (Figure~\\ref{figsbr}).  Secondly, we find that the majority ($64.9\\% \\pm 5.4\\%$  for  $M_\\star \\geq 5\\times 10^{10}$ $M_\\odot$, and $58.5\\% \\pm 7.7\\%$ for  $M_\\star \\geq 10^{11}$ $M_\\odot$)  of  massive galaxies  at $z=2-3$  have  low $n\\le$~2, while the corresponding fraction among massive systems at $z\\sim$~0 is five times lower. Most ($\\sim72\\%$) of these massive galaxies at $z=2-3$ with low $n\\leq2$ have $r_e > 2$ kpc, and therefore complement the ultra-compact galaxies. We further explore the meaning of a   S\\'ersic index $n\\le2$ at $z=2-3$, and present evidence that {\\it most of the  massive galaxies with  $n\\le2$ at $z=2-3$, particularly the extended ($r_e>2$ kpc) ones, likely host a prominent disk},   unlike the majority of massive galaxies at $z\\sim 0$. Our evidence is based on rest-frame optical morphologies, ellipticities, artificial  redshifting experiments,   as well as bulge-to-total ratios  from bulge-plus-disk  decompositions of extended systems.  \\vspace{3mm} {\\it  2.  Star formation rates:} We estimate star formation rates using IR luminosities (8-1000 $\\mu$m) derived from  the $Spitzer$ 24 $\\mu$m flux for massive GNS galaxies having  a secure MIPS 24~$\\mu$m counterpart and a 24~$\\mu$m  flux exceeding the 5$\\sigma$ detection limit of 30 $\\mu$Jy. AGN host candidates are excluded because the inferred IR luminosities overestimate the true star formation rate.  We find a strong link between galaxy structure and SFR. Among the non-AGN massive ($M_\\star \\geq 5\\times 10^{10}$ $M_\\odot$) galaxies   at $z=2-3$ with SFR$_{\\rm IR}$  high enough to yield a 5$\\sigma$  (30~$\\mu$Jy) $Spitzer$  24 $\\mu$m detection, the majority  (84.6 $\\pm 10.0$\\%)  have low $n\\leq 2$. While the  $n\\leq 2$ disky systems have a wide range of SFR$_{\\rm IR}$ (53 to 1466 $M_\\odot$ yr$^{-1}$ at $z=2-3$), they include the systems of the highest SFR$_{\\rm IR}$ at both $z=1-2$ and $z=2-3$. In contrast, the massive ultra-compact objects at $z=2-3$ are less likely by a factor of $\\sim 2.2$ to have SFR$_{\\rm IR}$ above the detection limit, compared to the whole sample of non-AGN massive galaxies.  \\vspace{3mm} {\\it 3. AGN activity: } Using a variety of techniques (X-ray properties, IR power-law, and IR-to-optical excess) to identify AGN, we find that  49/166 ($29.5 \\pm 3.5\\% $) of the massive galaxies at $z=1-3$ are AGN candidates. The AGN fraction rises with redshift, increasing from $17.9 \\pm 6.1\\%$ at $z=1-1.5$ to $40.3 \\pm 8.8 \\% $ at $z=2-3$ (Table~\\ref{tabagnfrak}).  We find a relationship between host galaxy structure and AGN activity that complements the relationship between SFR and structure. Among massive galaxies at $z=2-3$,  AGN appear to be found preferentially in galaxies that are not ultra-compact, as evidenced by the fact that most ($80.6 \\pm 7.9\\%$) AGN  hosts have $r_e > 2$ kpc. In fact, at $z=2-3$, the AGN fraction in ultra-compact galaxies is $\\sim2.7$ times lower than in extended galaxies ($20.0 \\pm 16.3\\%$ versus $53.2 \\pm 10.0\\%$).  {\\it Thus, ultra-compact galaxies  appear quiescent in terms of both  SFR and AGN activity.} In terms  of their S\\'ersic index $n$, a large fraction ($64.6 \\pm 10.7\\%$) of AGN hosts at $z=2-3$ have disky ($n\\leq2$) morphologies.  \\vspace{3mm} {\\it 4.  Cold gas content:} We apply a  standard Schmidt-Kennicutt law (Kennicutt 1998) to the SFR$_{\\rm IR}$ of the non-AGN host candidates. The high estimated SFR$_{\\rm IR}$  suggest that copious cold gas reservoirs are present. We estimate that the average cold gas surface density in non-AGN hosts ranges from $\\sim 136$ to  $\\sim 25,091$ $M_{\\odot}$~pc$^{-2}$ at $z=1-3$, with  a median value of $\\sim607$ $M_{\\odot}$~pc$^{-2}$ (Figure~\\ref{figfgas}). The implied cold gas fraction within the rest-frame optical half-light radius ranges from $6.5\\%$ to  $65.4\\%$, with a mean of $\\sim41\\%$ at $z=2-3$  (Figure~\\ref{figfgas}). The highest gas fractions at a given redshift are found among the less massive galaxies, consistent with downsizing.  \\vspace{3mm} {\\it 5.  Formation of  massive galaxies by $z=2-3$:} The massive galaxies at $z=2-3$ already have an average rest-frame optical  surface brightness within $r_e$  that can be up to 3-6 magnitudes brighter than $z\\sim0$ massive galaxies.  The associated high  stellar mass densities imply that massive galaxies at  $z=2-3$ must have formed via rapid, highly  dissipative events at $z>2$. Both gas-rich major mergers and  gas  accretion at $z>2$ are viable as their associated short  dynamical timescales and short gas cooling times at $z>2$ would  lead to a rapid buildup of mass. However, the large fraction ($\\sim65$\\%) of massive galaxies at $z=2-3$ with $n\\leq2$ and disky morphologies suggest that  cold-mode accretion at $z>2$  must have played an important role  in  the build-up of massive galaxies  by  $z=2-3$, since it may be challenging to  have  such a  large  fraction of of merger remnants with low $n \\le 2$   from isolated  gas-rich major mergers.  \\vspace{3mm} {\\it 6. Transformation of massive galaxies at  $z=2-3$ into present-day E and S0s:} In order for  massive galaxies at $z=2-3$ to evolve into   $z\\sim0$ massive systems (which are primarily E and S0s), they need to radically change their rest-frame optical structure and distributions of   ($n$, $r_e$). In particular they need to raise $n$  well above 2, increase $r_e$ by an average factor of 3-4, and dim the average rest-frame optical surface brightness. Dry major mergers can  induce  changes in galaxy size, S\\'ersic index, and stellar surface density,  but they may be too rare  to account for all the needed evolution.  Successive dry minor mergers have been shown to influence galaxy size, S\\'ersic index, and stellar surface density in  a similar direction. We suggest the transformation of massive $z=2-3$ galaxies into $z\\sim0$ galaxies will occur through a combination of dry major mergers, minor  mergers.  We will investigate in the relative importance and efficiency of these mechanisms in a future paper."
    },
    "1107/1107.0594_arXiv.txt": {
        "abstract": "The definition and measurement of magnetic reconnection in three-dimensional magnetic fields with multiple reconnection sites is a challenging problem, particularly in fields lacking null points. We propose a generalization of the familiar two-dimensional concept of a magnetic flux function to the case of a three-dimensional field connecting two planar boundaries. Using hyperbolic fixed points of the field line mapping, and their global stable and unstable manifolds, we define a unique flux partition of the magnetic field. This partition is more complicated than the corresponding (well-known) construction in a two-dimensional field, owing to the possibility of heteroclinic points and chaotic magnetic regions. Nevertheless, we show how the partition reconnection rate is readily measured with the generalized flux function. We relate our partition reconnection rate to the common definition of three-dimensional reconnection in terms of integrated parallel electric field. An analytical example demonstrates the theory, and shows how the flux partition responds to an isolated reconnection event. ",
        "introduction": "This paper presents a new method for measuring magnetic reconnection in a three-dimensional (3D) magnetic field. Reconnection is a fundamental physical process in any highly-conducting plasma, yet remains poorly understood owing to the challenging range of lengthscales involved \\citep{priest2000,birn2007}. In 3D magnetic fields, progress is hampered by the difficulty in defining and measuring reconnected flux, particularly if there are multiple interacting reconnection sites. It is this problem that we seek to address.  There are two contrasting ways to measure reconnection rates in 3D. The first uses the parallel electric field integrated along magnetic field lines\\citep{schindler1988}, while the second counts the transfer of flux between distinct flux domains. Here we pursue the second approach, where the task is twofold: to define a partition of the flux and to measuring the rate of transfer between fluxes in this partition. We call this a reconnection rate with respect to a partition, or a  \\emph{partition reconnection rate}, to distinguish it from the first case. Such a partition reconnection rate can capture only reconnection processes that change fluxes between the partition domains, and not those within any individual flux domain. However, this rate is in many applications the most relevant information, detemining the stability and  dynamics of the system.  In a two-dimensional (2D) field ${\\bf B}=B_x(x,y){\\bf e}_x + B_y(x,y){\\bf e}_y$, there is a natural choice for such a partition and correspondingly for the reconnection rate: write ${\\bf B}$ in terms of a \\emph{flux function} $A(x,y)$ where ${\\bf B}=\\nabla\\times A\\,{\\bf e}_z$. The different fluxes in the partition correspond to the regions of the plane bounded by separatrices, which are the global stable and unstable manifolds of hyperbolic nulls (x-points). The magnetic flux within each such region (per unit height in the ignorable $z$ direction) is measured by the difference in $A$ between appropriately chosen nulls. These fluxes are invariant under an ideal evolution, while in a non-ideal evolution the change in fluxes is measured exactly by the change in the values of $A$ at the (discrete) set of null points. This defines an unambiguous global reconnection rate which is readily computed even in turbulent 2D fields with many nulls \\citep{servidio2009}.  In 3D, the situation is more complicated. Here a natural partition also arises from the existence of null points in the domain. The 2D invariant manifolds (fan surfaces) associated with these null points form a coarse but natural  partition of the flux\\citep{longcope2009,haynes2010}. This inherent topological structure has been used successfully to quantify reconnection\\citep{pontin2011a}. Nulls and separators in particular are favoured locations to detect global changes in connectivity because they are locations where distinct flux domains come into close proximity.  Another possible flux partition arises in a field connected to a physical boundary, such as the photosphere of the Sun, where the sign of the normal field component divides the boundary into regions of positive and negative magnetic polarity. This boundary partition extends to a partition of magnetic flux connected to the boundary, by following the field lines from the different polarity regions into the volume\\citep{longcope2005}.  There are, however, many examples of magnetic fields where either the above described partitions are too coarse or null points do not exist in the domain. Examples are a single coronal loop or the magnetic field in a tokamak. A  generic 3D magnetic field in this situation does not possess a foliation of flux surfaces. This ``worst case\" is the situation which we want to investigate here.  More specifically, we assume a simply-connected domain in which  all field lines stretch between two planar boundaries. In magnetospheric reconnection studies  this is often referred to as the ``guide field'' case. We present a general way to define the flux partition in such a field, using distinguished hyperbolic orbits, and measure reconnection by introducing a generalized version of the 2D flux function $A$. Not only is this a natural topological flux partition when there are no magnetic null points, but it retains the simplicity of the 2D method when it comes to measuring the partition reconnection rate.  As may be expected, there are some complexities that do not arise in the 2D case. Chief among these is the possibility of chaos in the field line mapping. This is well-known from the study of area-preserving mappings as models for the magnetic field in toroidal fusion devices \\citep{morrison2000}. Indeed, recent results in tokamak experiments show heating on the vessel walls consistent with the breakdown of confinement and chaotic transport of magnetic flux through homoclinic tangles, as found in numerical simulations \\citep{evans2005,borgogno2008,viana2011}. Though they were recognised by Poincar\\'{e} in the 19th Century, it is only recently that detailed analysis of the structure of homoclinic tangles has been applied to measure and predict the transport of trajectories in these chaotic regions. A primary application has been 2D fluids with time-dependent velocity fields \\citep{romkedar1990a,ryrie1992,beigie1994,mancho2006}. Here, we show how these ideas can be applied to define and measure a natural reconnection rate in 3D magnetic fields. ",
        "conclusions": "\\label{sec:conclusion}  We have proposed a method to define and measure reconnection in a 3D magnetic field stretching between two boundaries. The flux is partitioned using the global manifolds of hyperbolic fixed points of the field line mapping between the boundaries. Individual fluxes in the partition are defined as differences between the values of a generalized flux function ${\\cal A}(x,y)$ at fixed points. This is a natural generalization of the flux function in a 2D magnetic field, and maintains the key advantage that the reconnection rate (with respect to this partition) is measured simply by the rate of change of ${\\cal A}$ at fixed points. The associated partition reconnection rate is unique, and straightforward to compute: the main computational effort required is identifying the fixed points in the field line mapping at successive times.  \\citet{petrisor2002} has previously recognised that reconnection may be related to values of an \\emph{action} function (essentially ${\\cal A}$) on hyperbolic orbits, though that analysis was limited to non-twist area-preserving maps. Our results are more general, and have been derived in a more physical way relating directly to the magnetic field itself. The interpretation of ${\\cal A}$ as an action integral brings out a deep connection with the magnetic structure. \\citet{cary1983} have shown that ${\\cal A}(x,y)$ is the action in a variational formulation leading to the equations of the magnetic field lines. In other words, given ${\\bf A}$, and assuming displacements $\\delta{\\bf x}$ with ${\\bf A}\\cdot\\delta{\\bf x}=0$ at the end-points, the Euler-Lagrange equations that extremise ${\\cal A}$ are the equations of the field lines.  Future work will consider in more detail how our partition reconnection rate compares with reconnection defined using the integrated parallel electric field, such as in numerical simulations (e.g. \\citet{pontin2011}). There are certainly differences: we have seen in Section \\ref{sec:example} how the visibility of a reconnection region (as defined with $E_{||}$) to our flux partition depends on its location. We interpret this as a difference in the \\emph{topological effectiveness} of the reconnection, from the point of view of the flux partition. One way to refine the partition would be to integrate  ${\\cal A}$ over more than one iteration of ${\\bf F}_1$, measuring the values of ${\\cal A}$ at the larger number of periodic points.  The example in Section \\ref{sec:added} also demonstrated that the flux in each direction across a partial barrier may be much larger than the net flux, and furthermore may change under a non-ideal evolution even if the net flux remains constant. Our existing flux partition is insensitive to such changes. However, we note that it is possible to measure the flux of an individual lobe with the formula \\begin{equation} \\Phi(L) = \\sum_{k=-\\infty}^{\\infty}\\left[{\\cal A}\\big({\\bf F}_1^k({\\bf p}_2)\\big) - {\\cal A}\\big({\\bf F}_1^k({\\bf p}_1)\\big) \\right], \\end{equation} where ${\\bf p}_2$ and ${\\bf p}_1$ are the defining pips of the lobe. This may be proved using a similar geometrical argument to Theorem \\ref{thm:netflux}, and is a straightforward generalization of a similar formula for lobe area in area-preserving maps \\citep{mackay1984,mackay1986,easton1991}. Taking account of chaotic regions in this way is likely to be of particular importance in fields where heteroclinic tangles fill large areas of the plane, with multiple intersections. In this case, the structure of the field is dominated by chaotic regions. While the flux partition may still be defined in the same way using ${\\cal A}$ values at the fixed points, the geometrical interpretation is less clear and requires further investigation.  Finally, a limitation of this work at present is our restriction to periodic fields where $B_z(x,y,1)=B_z(x,y,0)$. In future, we hope to relax this restriction thanks to recent theoretical progress in aperiodic dynamical systems \\citep{malhotra1998}. The notion of hyperbolic trajectories, and their stable and unstable manifolds, may be extended to this more general setting."
    },
    "1107/1107.0077_arXiv.txt": {
        "abstract": "Images recorded with MegaCam are used to investigate the recent ($t \\leq 0.25$ Gyr) star-forming history (SFH) of the Local Group Sc galaxy M33. The data sample the entire star-forming disk, as well as areas immediately to the north and south of the galaxy. The properties of the stellar disk change near R$_{GC} = 8$ kpc. Within this radius the luminosity function of main sequence stars indicates that the star formation rate (SFR) has been constant with time during at least the past 250 Myr, while at larger radii the SFR has declined during this same time period. That the recent SFR in the inner disk has been constant suggests that M33 has evolved in isolation for at least the past $\\sim 0.5$ Gyr, thereby providing a constraint on the timing of any recent interaction with M31. The color of the main sequence ridgeline changes with radius, suggesting a gradient in extinction of size $\\Delta A_V/\\Delta R_{GC} = -0.05$ magnitudes kpc$^{-1}$. The fractional contribution that young stars make to the total mass of the stellar disk changes with radius, peaking near 8 kpc. Evidence is also presented of systematic spatial variations in the SFH of the disk, such that the SFR during the past 100 Myr in the southern half of the galaxy has been $\\sim 0.4$ dex higher than in the northern half. Finally, structures with sizes spanning many kpc that contain blue objects -- presumably main sequence stars that formed during intermediate epochs -- are identified near the disk boundary. It is argued that these are tidal features that were pulled from the main body of M33 and -- in some cases -- are the fossil remnants of star formation that occured in an extended disk during intermediate epochs. ",
        "introduction": "The largest members of the Local Group have not evolved passively, but have experienced interactions that are consistent with hierarchal assembly. Perhaps the most overt signatures of this activity in the Local Group are the stellar streams that have been detected in the outer regions of M31 (e.g. Ibata et al. 2007; Tanaka et al. 2010). The location and morphology of the Sagittarious dwarf galaxy (Ibata et al. 1994), coupled with the presence of coherent streams in the Galactic disk (e.g. Helmi et al. 2006) and globular clusters that follow an age-metallicity relation that is distinct from that defined by the majority of Galactic globular clusters (e.g. Forbes \\& Bridges 2010), indicate that the Galaxy also has not evolved in isolation. There is evidence that the effects of recent interactions on the properties of the Galactic disk have not been as great as they have been on the M31 disk (e.g. Hammer et al. 2007).  As the nearest Sc galaxy, M33 is on a par with the LMC in serving as a benchmark for understanding the stellar content and evolution of late-type galaxies. However, during the past few years, evidence has been presented that M33 and M31 may have interacted during cosmologically recent epochs, and such an interaction may have affected the properties of the M33 disk. Braun \\& Thilker (2004) find an HI structure that appears to link M33 and M31, and Bekki (2008) models this as a tidal bridge that formed during an interaction 4 -- 8 Gyr in the past. Putman et al. (2009) investigate the possible orbital trajectories of M31 and M33 as constrained by their present-day locations and radial velocities. They conclude that at some point during the past $\\sim 1 - 3$ Gyr the tidal radius of M33 may have been $\\leq 15$ kpc due to an encounter with M31, which would have greatly affected the disk of M33. Simulations discussed by McConnachie et al. (2009) also support the notion of an interaction between M31 and M33 that may have disturbed the M33 disk. Their simulations suggest that the pericentric separation $\\sim 2.5$ Gyr in the past may have been as small as 40 kpc.  The spatial distribution and kinematic properties of the interstellar medium (ISM) of M33 show signatures of a recent tidal interaction. The HI in the outer regions of M33 has an S-shaped morphology and is distributed over $28 \\times 38$ kpc$^2$, with almost one fifth of the gas by mass located outside of the star-forming disk (Putman et al. 2009). While S-shaped structures are classical signatures of tidal interactions in simulations (e.g. Barnes \\& Hernquist 1992), the arms of the HI emission associated with M33 are in the opposite direction of disk rotation, possibly suggesting an encounter in the opposite direction of M33 disk rotation.  The HI in M33 also has a high velocity dispersion when compared with other late-type galaxies, with a mean $\\sigma = 18.5$ km sec$^{-1}$ (Putman et al. 2009). While high velocity dispersions in the outer regions of disks are not uncommon, and have been interpreted as the product of heating by halo sub-structures (e.g. Herrmann et al. 2009), the entire gas disk of M33 is kinematically hot. Block et al. (2007) argue that the warp in the ISM at 5 kpc could result from the accretion of gas with an angular momentum vector that was distinct from that of the main body of the disk, which might be expected to occur due to an interaction. Finally, while molecular clouds in M33 have some characteristics that are similar to their counterparts in the Milky-Way and M31 (Sheth et al. 2008), the ratio of atomic to molecular gas in M33 is high, and the fraction of molecular gas in giant molecular clouds is low when compared with the Galaxy (Verley et al. 2010). Both of these properties are consistent with a disturbed ISM.  \\subsection{Stars and Gas in M33: A Brief Review}  Stars and gas in M33 provide a fossil record of the events that shaped its past. Because of the obvious challenges presented by crowding, much of the previous work on the resolved stellar content of M33 has examined the outer regions of the disk. There are hints that the orbits of disk stars in M33 may have been disrupted, such that the radial age distribution differs from that in more isolated late-type spirals (Gogarten et al. 2010). In fact, it has been known for some time that stars with properties that are consistent with a disk origin appear at large radii in M33, as might be expected if the disk was disrupted. Davidge (2003) found that red giant branch (RGB) stars located between 14 and 17 kpc (7 -- 9 disk scale lengths) from the center of M33 have [M/H] $\\sim -1 \\pm 0.3$ dex. This metallicity is higher than expected for stars that belong to a classical metal-poor halo, but is within the range expected for disk objects. A similar situation occurs in M31, where moderately metal-rich stars are distributed throughout much of the outer regions of the galaxy (e.g. Bellazzini et al. 2003).  Davidge (2003), Block et al. (2007) and Verley et al. (2009) find evidence of a large population of C stars in the outer disk of M33. The discovery of a large number of C stars is suggestive of an elevated star formation rate (SFR) 1 -- 3 Gyr in the past, which is the likely time of an interaction with M31 (Putman et al. 2009). There are indications that these stars may belong to a population that is distinct from the inner disk population. Indeed, if -- as argued by Block et al (2007) -- the change that is seen in the $8\\mu$m profile is driven by the circumstellar dust contents of M giants and C stars, then this is indicative of a change in stellar content. Block et al. (2007) further suggest that the formation of the C stars may be tied to the events that caused the warp in the M33 HI disk at R$_{GC} = 5$ kpc.  Barker \\& Sarajedini (2008) explore the chemical evolution of the outer regions of M33. The chemical evolution deduced from CMD morphology does not follow that expected for a closed box; rather, a time-variable rate of infall is required to explain the enrichment history. Barker \\& Sarajedini (2008) find that the peak rate of gas accretion occured $3 - 7$ Gyr in the past, during which time 50 -- 60\\% of all gas that has fallen onto the disk was accreted. The rate of infall has since plunged, with only 10\\% of all infalling gas being accreted during the past 3 Gyr. While Barker \\& Sarajedini (2008) discuss these results in the context of protracted disk assembly, the drop in the infall rate some 3 Gyr in the past coincides with the timing of an interaction between M31 and M33 as predicted by both Putman et al. (2009) and Bekki (2008). A modest infall rate at the present day would be expected if circumgalactic gas was stripped by an encounter with another galaxy.  Past studies have found evidence for systematic radial variations in the star-forming history (SFH) of the M33 disk. Bastian et al. (2007) find that the number of young stellar groups drops suddenly at radii $> 4$ kpc, which they attribute to a change in the star-forming ISM, and it is only in the central 4 kpc that Park et al. (2009) find clusters with log(t) $\\leq 7.8$. The impact of systematic gradients in the SFH are also seen in integrated light. This is demonstrated by Munoz-Mateos et al. (2007), who find a radial gradient in the specific SFR (sSFR) of M33 based on measurements in the UV and IR.  Cioni et al. (2008) examine the distribution of M giants and C stars in M33, and find that the C/M ratio peaks $\\sim 35$ arcmin ($\\sim 10$ kpc) from the galaxy center. The radial variation in C/M is likely a product of gradients in both mean metallicity and age, in the sense that younger mean ages and higher metallicities occur at smaller radii. Using C/M as a metallicity indicator, Cioni (2009) finds a progressive radial decrease in metallicity with increasing radius in M33 out to 8 kpc, while at larger radii the metallicity profile is flat. Cioni (2009) interprets the break in the metallicity profile as an interface between inner and outer disk/halo components.  The spectral energy distribution of the M33 disk at visible wavelengths is consistent with insight-out formation (Li et al. 2004), although the radial age distribution of stars may be affected by processes other than the mechanics of disk assembly. Williams et al. (2009) use deep ACS images to examine stellar content over a range of radii in M33. They find that mean age drops as radius increases when R $< 8$ kpc, but at R $> 8$ kpc mean age increases with radius. Such a radial age inversion is consistent with simulations of disk evolution (e.g. Roskar et al. 2008; Sanchez-Blazquez et al. 2009), and can be understood in the context of inside-out disk formation coupled with the dynamical evolution of disk stars, which causes the outermost regions of the disk to be populated by stars that formed at smaller radii, but have migrated outwards. The most overt signature of such processes at work in M33 is the difference in scale lengths that characterize the distributions of young stars and old stars, in the sense that old stars are distributed over larger spatial scales than young stars (e.g. Verley et al. 2009; Davidge et al. 2011).  Metallicity gradients are a natural consequence of disk assembly in a hierarchal universe (e.g. Colavitti et al. 2009). However, if M31 passed close to M33 then tidal forces may have mixed gas throughout the disk, thereby flattening pre-existing gradients in the ISM. Rosolowsky \\& Simon (2008) examine the emission spectra of HII regions that span a range of radii in M33 and find that $\\Delta$[O/H]/$\\Delta$R $\\sim -0.027 \\pm 0.012$. This is substantially smaller than the mean [O/H] gradient in the 39 disk galaxies studied by Zaritsky et al. (1994), for which $<\\Delta$[O/H]/$\\Delta$R$> = -0.058 \\pm 0.009$, where the uncertainty is the standard error in the mean. A caveat is that giant HII regions, upon which many of the measurements in the Zaritsky et al. (1994) compilation are based, may define steeper gradients than those computed from all HII regions, regardless of size (Magrini et al. 2010).  Rosolowsky \\& Simon (2008) find considerable scatter about the mean relation between [O/H] and radius in M33, which they attribute to local inhomogeneities in the ISM. An intriguing result is that HII regions with low [O/H] (and hence presumably low total metallicity) are found at small radii, suggesting that material that has undergone only modest amounts of chemical enrichment is present in the part of M33 that is expected to be the most chemically mature. The presence of such material could be due to the tidally-induced mixing of gas from the outer regions into the central regions of the galaxy. Alternatively, Magrini et al. (2010) argue that the comparatively low mean metallicity computed for HII regions in the central kpc of M33 may be a result of selection effects in surveys that require the detection of [OIII]4363, which is a diagnostic of electron temperature. Still, it is evident from their Figure 19 that the mean [O/H] in the central kpc remains lower than that at 1.5 kpc even if HII regions that do not have [OIII]4363 detections are included in the sample.  Very young stars might be expected to track the metallicity trends defined by HII regions. However, blue supergiants (BSGs) in M33 define a steeper radial metallicity gradient than HII regions (U et al. 2009), with a slope that is consistent with the mean in the Zaritsky et al. (1994) HII sample. A possible explanation for this difference in gradients is that U et al. (2009) do not find stellar counterparts to the oxygen-poor HII regions detected at small radii by Rosolowsky \\& Simon (2008); this may be a consequence of the modest number of BSGs studied.  The ages of HII regions and PNe differ by at least a few hundred Myrs, and so insights into the recent chemical enrichment history of the M33 disk can be obtained by comparing abundance gradients deduced from these types of nebulae. Magrini et al. (2009) find that $\\Delta$[O/H]/$\\Delta$R = $-0.031 \\pm 0.013$ amongst M33 PNe, and this agrees with the gradient measured from HII regions by Rosolowsky \\& Simon (2008). Metallicity gradients are expected to steepen with time if there is infall, and Magrini et al. (2009) attribute the lack of evolution of the [O/H] gradient to a low accretion rate during the past few Gyr. Such a drop in infall rate during intermediate epochs is in broad agreement with the conclusion reached by Barker \\& Sarajedini (2008) from an investigation of CMDs. Alternatively, a static abundance gradient with time could also signal the accretion of gas with low or modest amounts of enrichment, at a fortuitous rate that is sufficient to balance any increase in metallicity that may arise from the recycling of nuclear processed material into the ISM.  Some properties of M33 may not be typical of late-type spiral galaxies in general. For example, when compared with other late-type galaxies, the star-forming activity in M33 appears to be subdued. Indeed, while the sSFR of M33 is close to the midpoint of spiral galaxies in general, it is near the low end of galaxies of similar morphological type (Munoz-Mateos et al. 2007). The radial sSFR gradient in M33 is also steeper than in the majority of galaxies (Munoz-Mateos et al. 2007). This may indicate a low sSFR at small radii, although the circumnuclear stellar content of M33 is similar to that in other nearby late-type spirals (Davidge 2000; Davidge \\& Courteau 2002).  The gas within 8 kpc of the center of M33 appears to have a sub-critical density for triggering star formation (Martin \\& Kennicutt 2001), and this is one possible explanation for the low sSFR when compared with other late-type spirals. Considering the low gas density, the chemical properties of M33 are suggestive of a star-forming efficiency (SFE) that is higher than that among other nearby galaxies, and the SFE increases to progressively larger radii (Magrini et al. 2010). That star-forming activity occurs on a large scale in an environment with a low density may be related to the shear rate (Martin \\& Kennicutt 2001), which influences the frequency of collisions between molecular clouds, or an increase in ISM pressure caused by material falling into the galaxy, coupled with a kinematically hot interstellar environment (Putman et al 2009). The gravitational field of stars in the disk may also facilitate the collapse of star-forming material (e.g. Verley et al. 2010).  \\subsection{The Current Study}  The youngest stars in late-type galaxies play a key role in defining the classical morphological characteristics of these systems, and if M33 has experienced a recent interaction then residual signatures of such an event may be evident in the spatial distribution of these objects. Massive young stars are also brighter than the vast majority of older stars, and so are less prone to crowding in ground-based observations. In the present paper, $u*$ and $g'$ images obtained with MegaCam on the 3.6 metre Canada-France-Hawaii Telescope (CFHT) are used to investigate the properties of young stars in M33.  Two specific aspects of young stars in M33 are examined in this paper. First, we investigate the recent SFH of the M33 disk, with emphasis on how it has changed with time and location during the past few 100 Myr. There are considerable uncertainties associated with SFHs that are based on integrated light indicators, and these uncertainties are avoided by employing star counts as a probe of SFH. While area-to-area variations in SFR may occur within a galaxy due to the passage of spiral density waves or stochastic effects, when averaged over long timescales and large areas the SFRs of isolated star-forming disks are expected to be more-or-less constant with time. Aside from dense galaxy clusters, where there may be a high-density ambient intergalactic medium, galaxy-wide departures from a constant SFR in intermediate mass galaxies are likely triggered by tidal encounters.  Second, we search for young or intermediate age stars outside of the main body of the M33 disk. One might expect a mixed bag of such objects at large distances from the galaxy center. The outermost regions of the stellar disk may contain stars that migrated from smaller radii (e.g. Sellwood \\& Binney 2002), while the passage of spiral density waves might generate isolated areas of star formation (Bush et al. 2010). Tidal effects may pull gas from the inner regions of disks, and studies of interacting galaxies indicate that star formation is not uncommon in tidal debris fields (e.g. Neff et al. 2005). Tidal interactions can also re-distribute existing stars over large areas (e.g. Teyssier et al. 2009). If a tidal plume formed due to an interaction between M31 and M33 then young/intermediate age stars, that either were present in the disk prior to the interaction or formed during the interaction, might be found in the circumgalactic environment.  An important aspect of this study is the near-complete spatial coverage of the M33 disk. The only areas missed in the stellar disk are those that fall in the gaps between detector banks. While crowding is an issue within 2 kpc of the galaxy center, a large fraction of the brightest blue stars are still resolved there with these data (\\S 6). Such complete areal coverage facilitates the suppression of stochastic variations in the SFH that may occur over kpc spatial scales. In this sense, the SFH deduced from our data is more akin to what might be derived from -- say -- fiber spectroscopic studies of distant galaxies, rather than from pencil-beam surveys of distinct areas within nearby galaxies that typically cover only a few kpc$^2$.  Observations in $u*$ are another important aspect of this program. The $u*$ filter is well-suited to investigations of intrinsically luminous main sequence stars, not only because they have effective temperatures that places the peak of their spectral energy distributions shortward of $0.4\\mu$m, but also because of the enhanced contrast with respect to the (predominantly red) unresolved stellar body of M33 when compared with filters having longer central wavelengths. Observations in $u*$ are also beneficial when searching for diffuse ensembles of young and intermediate age stars in the outermost regions of the disk, where the fractional contamination from unresolved background galaxies, the majority of which have red colors (e.g. Davidge 2008; Mouhcine \\& Ibata 2009), is substantial. An obvious downside to observing in $u*$ is that there is increased sensitivity to reddening variations when compared with filters that have longer effective wavelengths.  The distance modulus computed by Bonanos et al. (2006), $\\mu_0 = 24.92$, is adopted for this study. This distance estimate is based on a geometric calibration (the orbital properties of stars in an eclipsing binary), and is in excellent agreement with recent distances computed from the properties of B supergiants and ACS-based observations of the RGB-tip (U et al. 2009). Galleti et al. (2004) compare various distance moduli of M33, and the results are summarized in their Figure 10. While the Bonanos et al. (2006) value is near the upper end of the distance estimates considered by Galleti et al. (2004), it still falls well within the scatter envelope. In any event, the basic results of this study will not change greatly if a distance modulus that is -- say -- 0.2 mag smaller were to be adopted.  The line-of-sight extinction is assumed to originate in a uniform absorbing sheet in front of M33, and the adopted total extinction is the result of combining estimates of the foreground and internal components. The foreground component is taken from Schlegel et al. (1998), for which A$_B = 0.18$, while the internal component is that computed for M33 by Pierce \\& Tully (1992), for which A$_B = 0.16$; thus, the total extinction is A$_B = 0.34$ magnitudes. In \\S 4 it is demonstrated that this model correctly predicts the mean extinction towards main sequence stars in M33. Still, it should be kept in mind that dust in M33 is not uniformly distributed. U et al. (2009) find that line-of-sight dust extinctions among bright A and B supergiants vary by $\\sim 0.16$ magnitude in $(B-V)$, with a mean E(B--V) = 0.08. This mean is consistent with the combined foreground and internal reddening values used here. It is demonstrated in \\S 4 that reddening variations, as gauged by the width of the main sequence, are not substantial, with $\\Delta E(B-V) \\leq \\pm 0.1 - 0.2$ throughout much of the M33 disk.  The paper is structured as follows. Details of the observations and the photometric measurements are discussed in \\S 2 and \\S 3. The color-magnitude diagrams (CMDs) of stars and an assessment of reddening variations throughout the disk are presented in \\S 4. The recent SFH throughout the disk is investigated in \\S 5 by comparing the luminosity function (LF) of main sequence stars with models. The spatial distribution of young and intermediate age stars throughout the disk and its immediate surroundings is examined in \\S 6. A summary and discussion of the results follows in \\S 7. ",
        "conclusions": "Local Group galaxies are fundamental calibrators for stellar content investigations of more distant systems, many of which are late-type spiral galaxies. As the nearest Sc galaxy, M33 is thus an obvious target for understanding the pedigrees of more distant galaxies. In this study, moderately deep wide-field images recorded with the CFHT MegaCam have been used to probe the recent star-forming history of the M33 disk and its immediate surroundings. Main sequence stars that formed within the past few hundred Myr are resolved throughout a significant fraction of the $\\sim 4.5$ degrees$^2$ that we examine.  Our study yields three main results. First, the LF of main sequence stars indicates that -- when averaged over large parts of the disk -- the SFR within the central 8 kpc of M33 has not changed with time during the past few hundred Myr. This in turn suggests that the disk has not been disturbed for at least this period of time, given that tidal interactions will affect the SFR by -- for example -- re-distributing star-forming material.  Second, the LF of main sequence stars at distances in excess of 8 kpc from the galaxy center is not consistent with a constant SFR during recent epochs. In fact, the peripheral regions of M33 are not locations of star formation at the present day, and the projected density of very young main sequence stars shows a steep drop near R$_{GC} = 8$ kpc. This radial truncation in the young star-forming disk is a robust result, that is also seen in the distribution of UV light (Thilker et al. 2005).  Third, we have found that the outer regions of M33 harbor diffuse collections of blue objects with photometric properties that are consistent with those of main sequence stars with ages $\\sim 100$ Myr. Some of these structures appear to emerge from the disk, while others may be isolated. These structures may be tidal in nature, and/or may be spatially extended areas where star formation occured {\\it in situ}. The latter interpretation opens the possibility that while star formation at the present day is restricted to within $\\sim 8$ kpc of the galaxy center, during intermediate epochs it may have occured over a much larger region. The appearance of M33 at this time would have been very different from today, especially in the UV.  \\subsection{Evidence for an Interaction in the SFH, and the Long-Term Impact on M33}  A centrally concentrated upswing in star-forming activity is one consequence of a tidal encounter with another galaxy. Tidal torques remove angular momentum from gas in the disk, causing it to migrate inwards. Elevated levels of star formation in the inner regions of the galaxy are then triggered as the gas density builds and cools. Pre-existing stars in the disk gain angular momentum via interactions with the infalling gas clouds, and move to larger radii, producing an excess of old and intermediate age stars in the outer disk (e.g. Younger et al. 2008). If there was a recent interaction with M31 then elevated levels of central star formation may have occured in M33 -- is there evidence of such activity in the resolved stellar content?  If the disk of M33 is recovering from an uptick in star-forming activity that occured within the past few hundred Myr, which is the typical timescale for a starburst, then the SFR should decrease as one moves towards more recent epochs. Such a decrease in the SFR is expected as the gas supply is depleted by star formation and/or feedback. Surveys of M33 star clusters suggest that a large peak in star-forming activity may have occured $\\sim 0.1$ Gyr in the past (Chandar et al. 1999). However, the excellent match between the observed LFs and constant SFR models within 8 kpc of the galaxy center indicates that the global SFR has not changed during the past $\\sim 0.2$ Gyr. In fact, given that interactions fuel elevated levels of star formation over roughly a disk rotation time ($\\sim 0.3$ Gyr), then the results presented in this paper suggest that an interaction with M31 could not have occured during the past $\\sim 0.2 + 0.3 = 0.5$ Gyr.  Studies that examine older stellar generations provide additional constraints as to when an interaction between M31 and M33 may have occured. Williams et al. (2009) find that the SFR near the center of M33 peaked during the past 0.5 -- 2.5 Gyr, at which time the SFR jumped by an order of magnitude above the level that prevailed 2.5 -- 10 Gyr in the past. The time at which this jump in the SFR occured falls within the constraints set by the present-day positions and motions of M31 and M33 (Putman et al. 2009). Williams et al. (2009) also find that (1) the SFR varied with radius, as expected if gas is displaced by an interaction, and (2) the relative amplitude of the increase in the SFR during this same time interval drops with increasing R$_{GC}$, also as expected in the context of inward radial gas flows.  The global SFR during the past few hundred Myr has been lower than the long-term average (Williams et al. 2009), and the sSFR of M33 is lower than in most other late-type galaxies (Munoz-Mateos et al. 2007). These results are consistent with the depletion of cool gas throughout M33, as might occur if there was sustained large-scale star formation during intermediate epochs. The substantial jump in the central SFR during intermediate epochs and the subsequent drop in the SFR since that time suggests that M33 may have experienced a modest starburst; however, any starburst activity was not so dramatic as to present-day restrict star-forming activity to the central few kpc, as is seen in M82, where a large fraction of the outer stellar disk has been denuded of gas.  The drop in the infall rate 3 Gyr in the past infered from the chemical enrichment properties of the outer disk of M33 by Barker \\& Sarajedini (2008) may be understood in the context of the stripping of the circumgalactic M33 HI reservoir by M31. The drop in infall occured at a time that is consistent with the interaction timescales estimated by Putman et al. (2009). The disruption of such a circumgalactic reservoir will affect the overall appearance of M33 in two ways.  First, the removal of gas from the outer halo of a galaxy may affect the size of the star-forming disk. Models of ram pressure stripping indicate that star-forming disks contract as the circumgalactic gas reservoir that fuels star-forming material is removed (e.g. Bekki 2009). Of course, the Local Group environment is such that M33 has not been subjected to significant ram pressure stripping. Still, if the M33 gas reservoir was disrupted tidally then the net impact on the star-forming disk may be similar to that produced by ram pressure stripping: the appearance of a compact star-forming disk for an extended period of time (presumably a significant fraction of the Hubble time) until an external gas reservoir is restored. The sharp outer disk boundary at 8 kpc in M33 may be a consequence of a truncated angular momentum content in the circumgalactic gas reservoir.  Second, the removal of an external reservoir that replenishes disk gas will change the SFR, and hence the appearance of the galaxy. If the SFR during the past few hundred Myr has been 0.7 M$_{\\odot}$ year$^{-1}$ (Blitz \\& Rosolowsky 2006), then $0.7 \\times 0.2 \\times 10^9 = 1.4 \\times 10^8$ M$_{\\odot}$ of stars have formed in the past 0.2 Gyr. Adopting a 5\\% star formation efficiency (Inoue et al. 2004), then this level of star formation requires $2.8 \\times 10^9$ M$_{\\odot}$ of gas. This is sustantially lower than the present-day HI mass of M33 (Putman et al. 2009). To be sure, gas is recycled throughout the disk; still, the point to be made is that M33 is not gas-rich and the SFR, which has been constant for the past few hundred Myr, will gradually drop as the disk gas supply is depleted. The drop in SFR during the past 10 Myr that is predicted by the star counts in \\S 4.2 may signal the start of this process.  \\subsection{The Distribution of Stars and Gas in M33 as a Probe of Interaction Timing}  The radial distributions of gas and stars provide clues into the evolution of disks. Gas and stars are expected to have different distributions in isolated disks. This is due in part to the disk assembly process, as the angular momentum distribution of accreted gas will differ from that of stars that are already present. However, even in the absence of gas accretion, the distinct dynamical behaviours of stars and gas results in different radial distributions. Because of their small cross-sections when compared with typical stellar separations, stars form collisionless systems, while the relatively large cross-sections of gas clouds make them collisional systems. Stars may also evolve into extraplanar systems that are distinct from the gas distribution, even in isolated galaxies. Indeed, some fraction of disk stars at small radii may evolve into a thick disk, as a natural consequence of inside-out disk formation (e.g. Schonrich \\& Binney 2009).  The distributions of gas and stars in disks have been modelled, and the results can be compared with the observed distributions in M33. Here, we consider the cosmologically-based galaxy simulations examined by Martinez-Serrano et al. (2009). A break in the model light profile typically occurs at $\\sim 3$ disk scale lengths. The gas distribution in these models is roughly flat within the break radius, and gradually decreases beyond this point. Star formation continues past the break radius, occuring out to 4.3 scale lengths, at which point it ceases due to the low density of star-forming material. The ratio of gas mass to stellar mass peaks at the break radius, making this the point in the disk with the lowest photometrically-weighted age.  The HI density profile throughout much of the M33 disk is relatively flat (Corbelli 2003; Putman et al. 2009), in agreement with the Martinez-Serrano et al. (2009) models. The general trend of increasing specific SFR with radius in the main body of the disk predicted by the models is also consistent with the relative contributions made by young and old stars shown in Figure 15. Still, the muted evidence for star formation outside of the break radius in M33 is not consistent with the models, and the presence of tidal features in the HI distribution outside of the disk (Putman et al. 2009) points to an event that undoubtedly affected the evolution of M33. The presence of large-scale tidal features aside, the distributions of gas and stars within the break radius of M33, where the characteristic timescale for relaxation is shorter than at large radii, are at least qualitatively consistent with predictions made for disk galaxies evolving in the CDM cosmological paradigm. That the relative distributions of gas and stars in the main disk of M33 are consistent with that of an isolated disk suggests that sufficient time has passed (i.e. a few disk rotations) to allow the distributions of stars and gas within the disk to regain at least the semblance of an equilibrium state.  \\subsection{The M31-M33 Connection}  An interaction between M31 and M33 would almost certainly have imprinted signatures in the stellar content of the former. McConnachie et al. (2009) suggest that an interaction with M33 would have disturbed the outer disk of M31, causing a warp and displacing disk stars to large radii. An interaction with M33 may also have triggered an inward flow of gas in M31, depleting gas reserves at large radii. Indeed, the gas content of M31 at large radii is deficient when compared with models of isolated evolution (Yin et al. 2009). The interaction may also have produced elevated SFRs in M31, and a sharp drop in overall star-forming activity would occur as gas is consumed during star formation or is disrupted by feedback. Williams (2002) concludes that a global downturn in star-forming activity occured in M31 $\\sim 1$ Gyr in the past. Given that starbursts appear to last no more than a Gyr (e.g. Leitherer 2001), then this result -- if due to an interaction with M33 -- suggests that M31 and M33 interacted no more than $\\sim 2$ Gyr ago, which is consistent with the age ranges estimated by Putman et al. (2009) and McConnachie et al. (2009). Of course, all of these properties of M31 could also be explained by an interaction with another (possibly now defunct) companion.  If there was a close interaction between M31 and M33 then some of the features traditionally associated with M31 companions could instead be remnants of this interaction. Consider the M31 Giant Southern Stream (GSS; Ibata et al. 2001). The simulations run by Bekki (2008) indicate that M33 may leave a gaseous debris stream near M31, and Fardal et al. (2008) conclude that the GSS may have originated from a spiral galaxy. Indeed, the GSS contains stars with a metallicity as high as [Fe/H] $\\sim -0.5$ (Guhathakurta et al. 2006), which is consistent with an origin in the M33 disk. The velocity dispersion of stars in the GSS (e.g. Ibata et al. 2004) is also consistent with an origin in a kinematically cold environment, such as a disk. Finally, the GSS is 85 kpc behind M31 (e.g. Tanaka et al. 2010), and simulations suggest that the GSS was produced by an object on a highly eccentric orbit (Font et al. 2006; Fardal et al. 2006; 2007), such as M33.  While some properties of the GSS can be rationalized in the context of an interaction with M33, efforts to link the GSS with such an event run into timing issues. The GSS has a dynamical age of a few hundred Myr (e.g. Font et al. 2006), which is more recent than the minimum probable age for an interaction estimated by Putman et al. (2009), and is not consistent with the time interval during which there has been a constant SFR throughout the M33 disk. Another issue is that the vast majority of stars in the GSS have an age $\\geq 4$ Gyr (Brown et al. 2006), and so fall outside of the nominal age range predicted for an interaction. Still, a modest population of young stars is present in the GSS (e.g. Figure 17 of Brown et al. 2006). Recently formed stars in extraplanar environments may not be well-mixed with underlying stellar material, as they formed in gas concentrations, whereas pre-existing stars will tend to dissipate spatially. If this is the case in M31 then future observations may find concentrations of intermediate age stars in different fields than those sampled by Brown et al. (2006).  \\parindent=0.0cm  \\noindent  \\clearpage  \\begin{table*} \\begin{center} \\begin{tabular}{lll} \\tableline\\tableline Program ID & $u*$ & $g'$ \\\\ \\tableline 2003BF15 & 715153--156 & 714745 \\\\ & & 715150, 797, 935, 938 \\\\ & & 716134 \\\\ & & 718659, 894 \\\\ & & 719309, 585 \\\\ & & 724409 \\\\ & & 727548 \\\\ 2004BH06 & 774552 & \\\\ & 774554--557 & \\\\ 2004BH20 & & 757867, 868 \\\\ & & 762372, 373, 374, 378, 380, 381, 382, 466--472 \\\\ 2004BF26 & & 759110 \\\\ & & 761403, 560 \\\\ & & 765027, 128, 927 \\\\ & & 767646, 804, 809, 810, 975 \\\\ & & 769086, 541 \\\\ & & 773381 \\\\ \\tableline \\end{tabular} \\end{center} \\caption{Archival Data Used to Construct the Center Field Images} \\end{table*}  \\clearpage"
    },
    "1107/1107.1840_arXiv.txt": {
        "abstract": "We show that in the single component situation all perturbation variables in the comoving gauge are conformally invariant to all perturbation orders. Generally we identify a special time slicing, the uniform-conformal transformation slicing, where all perturbations are again conformally invariant to all perturbation orders. We apply this result to the $\\delta{N}$ formalism, and show its conformal invariance. ",
        "introduction": "A large class of extensions of general relativity is described in the context of scalar-tensor theories of gravity~\\cite{STreview}, where a space-time metric $g_{\\mu\\nu}$ is coupled to a scalar field $\\phi$, with the other matter contents such as fermion fields being minimally coupled to gravity. By an appropriate conformal transformation \\begin{equation}\\label{ct} g_{\\mu\\nu} \\to \\overline{g}_{\\mu\\nu} = \\Omega^2g_{\\mu\\nu} \\, , \\end{equation} we can move to another frame. A prime example is the Einstein frame, where after the conformal transformation (\\ref{ct}) the gravitational action becomes the Einstein-Hilbert one. These conformal frames are mathematically equivalent, so one can work in any frame as long as mathematical manipulations are concerned. It is, however, not explicitly clear whether these conformally related frames also enjoy physical equivalence~\\cite{Deruelle:2010ht}.   In cosmology, we encounter various frames of the metric which are related by conformal transformation, such as Jordan frame, Einstein frame, string frame and so on. The physical equivalence between these frames is especially important, since the cosmological observations with improving sensitivity can probe gravitational theories on the largest observable scales~\\cite{Song:2008vm}. Further, some models of the early universe which accounts for the primordial perturbations incorporate non-minimal coupling between gravity and a scalar field, such as the recently proposed standard model Higgs inflation~\\cite{Bezrukov:2007ep}. Usually one moves from the original, Jordan frame to the Einstein frame by a conformal transformation, and then computes perturbations there. The invariance of the vector and tensor perturbations, and that of the scalar perturbation in the comoving gauge, are shown up to second order~\\cite{equivalence}, but it is not clear if this invariance holds fully non-perturbatively. Regarding precise upcoming experiments that may detect non-linear signatures such as non-Gaussianity, it is important to clarify the full non-linear invariance of the perturbations~\\cite{Chiba:2008ia}.   In this article, we study the conformal invariance of the cosmological perturbation. Our particular attention will be given to the curvature perturbation on super-horizon scales. We will show that in the single component case all perturbations in the comoving slicing are conformally invariant to fully non-linear order. Once this is given, we can see that in the context of the $\\delta{N}$ formalism it is a matter of gauge transformation between different coordinate systems which impose different slicing conditions.   As a temporal gauge condition, we take the conformal transformation factor $\\Omega$ to be unperturbed to full non-linear order. This may be called as the uniform-conformal-transformation slicing, or, uniform $\\Omega$ slicing (U$\\Omega$S). It is convenient to decompose $\\Omega$ into the background and perturbation as $\\Omega = \\Omega_0e^\\omega$, then the U$\\Omega$S means $\\omega=0$ as the slicing condition. Thus we have \\begin{equation} \\Omega = \\Omega_0(t) \\, . \\end{equation} Under this slicing condition it is obvious, almost a tautology, that all perturbed quantities are naturally invariant under the conformation transformation. This result applies even in multiple component situation: see (\\ref{UOmegaG-fR}) and (\\ref{UOmegaG-Fphi}) for relations implied by this slicing condition to non-linear orders of perturbation. Note that we may consider the action given by \\begin{equation} S=\\int d^4x\\sqrt{-g}\\left[{1 \\over 2} f (\\phi,R) -\\frac{1}{2} \\omega (\\phi)g^{\\mu\\nu} \\phi_{,\\mu}\\phi_{,\\nu}-V(\\phi)\\right]\\, . \\end{equation} Here, $\\Omega$ is a function of either the field $\\phi$ for $f=F(\\phi) R$-type gravity or a function of $F$ or $R$ for $f(R)$-type gravity, where \\begin{equation} F\\equiv \\frac{\\partial f}{\\partial R}\\,. \\label{defF} \\end{equation} In these cases, the U$\\Omega$S corresponds to the uniform-field slicing or the uniform-$F$ slicing, respectively, and we may call it as the UFS. ",
        "conclusions": "To summarize, we have studied the non-perturbative conformal invariance of the cosmological perturbations. When the universe is dominated by a single component, the U$\\Omega$S and the comoving slicing coincide and the comoving curvature perturbation $\\calR_c$ is conformally invariant fully non-perturbatively. Consistent with the conformal invariance of $\\calR_c$, we have shown that the $\\delta{N}$ formalism gives identical results irrespective of the choice of conformal frames.   When the universe is dominated by a multiple matter or field components, the comoving curvature perturbation is no longer conformally invariant. However, at and after the universe has settled down to a unique evolutionary trajectory, i.e. in an era when there remains only a single adiabatic mode, we recover the conservation of the comoving curvature perturbation on super-horizon scales and so it is conformally invariant.     \\subsection*{Acknowledgements}   We acknowledge the workshop ``Cosmological Perturbation and Cosmic Microwave Background'' (YITP-T-10-05) at the Yukawa Institute for Theoretical Physics, Kyoto University, where the main part of this work was done, and the workshop ``WKYC 2011 -- Future of Large Scale Structure Formation'' at Korea Institute for Advanced Study, where this work was initiated and completed. This work was supported in part by a Korean-CERN fellowship, Korea Institute for Advanced Study under the KIAS Scholar program, the Grant-in-Aid for the Global COE Program at Kyoto University, ``The Next Generation of Physics, Spun from Universality and Emergence'' from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan, JSPS Grant-in-Aid for Scientific Research (A) No.~21244033, JSPS Grant-in-Aid for Creative Scientific Research No.~19GS0219, and KRF Grant funded by the Korean Government (KRF-2008-341-C00022).     \\appendix \\renewcommand{\\theequation}{\\thesection.\\arabic{equation}}"
    },
    "1107/1107.1592_arXiv.txt": {
        "abstract": "{{Many of the high mass X-ray binaries (HMXRBs) discovered in recent years in our Galaxy are characterized by a high absorption, most likely intrinsic to the system, that can impede their detection at the softest X-ray energies. Exploiting the good coverage  obtained with sensitive  \\XMM\\ observations, we have undertaken a search for highly absorbed X-ray sources in the Small Magellanic Cloud (SMC), which is known to contain a large number of HMXRBs. After a systematic analysis of 62 \\XMM\\ SMC observations, we obtained a sample of 30 sources with evidence of an equivalent hydrogen column density larger than 3$\\times$$10^{23}$ cm$^{-2}$.} {Five of these sources are clearly identified as HMXRBs, four being previously known (including three X-ray pulsars) and one, XMMU J005605.8--720012,  being reported here for the first time. For the latter, we present optical spectroscopy confirming the association with a Be star in the SMC. The other sources in our sample have optical counterparts fainter than magnitude $\\sim$16 in the V band, and many have possible NIR counterparts consistent with highly reddened early-type stars in the SMC.} {While their number is broadly consistent with the expected population of background highly absorbed active galactic nuclei, a few of them could be HMXRBs in which an early-type companion is severely reddened by local material.} ",
        "introduction": "Many observations of the Small Magellanic Cloud (\\SMC) in the X-ray energy band have led to the discovery of a large number of High Mass X-ray Binaries (HMXRBs). The number of  known HMXRBs in the \\SMC\\ (about one hundred, \\citet{Liu2005}) is much larger than expected by scaling the number of these sources seen in the Milky Way according to the mass ratio of the two galaxies (M$_{MW}$/M$_{SMC}$ $\\sim$50). This has been interpreted as evidence of a recent episode of star formation in the \\SMC\\ (\\citet{Majid2004}, \\citet{Shtykovskiy2007}, \\citet{Antoniou2010}). It is remarkable  that the \\SMC\\ contains only one supergiant system, SMC X-1. All the other HMXRBs in the \\SMC\\ consist of neutron stars accreting from Be type companions. Be X-ray binaries constitute also the largest class of HMXRBs in our Galaxy, but the relative number of supergiant systems is much higher than in the \\SMC, since 30 of the 114 Galactic HMXRBs  \\citep{Liu2006} are confirmed or suspected supergiant systems.  Many new Galactic HMXRBs have been discovered in the past few years with the INTEGRAL satellite, thanks to high sensitivity in the hard X-ray range, coupled with an extensive monitoring of the Galactic plane. These observations led to the recognition of the new class of Supergiant Fast X-ray Transients (SFXT) and to the discovery of persistent supergiant systems characterized by a high absorption (equivalent column density above a few 10$^{23}$ cm$^{-2}$), which escaped an earlier discovery because the high absorption severely suppresses their flux  below 10 keV (see, e.g., \\citet{Sidoli2010}, for a recent review).  To search for highly absorbed HMXRBs in the \\SMC\\, we carried out a dedicated analysis of  the \\XMM\\ data  collected during  the large program for the \\SMC\\ survey performed by \\citet{Haberl2008} and other observations of SMC targets. In Sect.~\\ref{sec:2}, we describe the X-ray observations and data processing and in Sect.~\\ref{sec:3} we describe the source detection and selection criteria that led us to identify a sample of 30 highly absorbed sources. Five of them are confirmed  HMXRBs, including four that were already known as X-ray binaries (Sect.~\\ref{sec:4}). In Sect.~\\ref{sec:5} we present optical spectroscopy of the new source, demonstrating that it is a Be system in the \\SMC. We finally discuss our results in Sect.~\\ref{sec:6}. ",
        "conclusions": "\\label{sec:7}   Motivated by the presence of several highly absorbed HMXRBs in our Galaxy, we have carried out a search for similar systems in the \\SMC\\, exploiting the nearly complete coverage of this galaxy obtained with \\XMM.  Our selection criteria, corresponding to an absorption threshold of $\\sim3\\times10^{23}$ cm$^{-2}$, were met by four known HMXRBs and led to the discovery of a new highly absorbed Be binary in the SMC. We also selected other 25 additional sources, among which other highly absorbed HMXRBs might be present. Since we also expect a good fraction of these sources to be AGNs, we can say that the SMC does not contain much more than $\\sim$10 persistent HMXRBs with intrinsic N$_H>3\\times10^{23}$ cm$^{-2}$. We cannot exclude the presence of a much larger population of intrinsically absorbed binaries of transient nature, which could not be detected in our survey if their quiescent luminosity is below a few 10$^{34}$ erg  s$^{-1}$."
    },
    "1107/1107.3418_arXiv.txt": {
        "abstract": "{We present a comprehensive study on the impact of the environment of compact galaxy groups on the evolution of their members using a multi-wavelength analysis, from the UV to the infrared, for a sample of  32 Hickson compact groups (HCGs) containing 135 galaxies. Fitting the SEDs of all galaxies with the state-of-the-art model of da Cunha (2008) we can accurately calculate their mass, SFR, and extinction, as well as estimate their infrared luminosity and dust content. We compare our findings with  samples of field galaxies, early-stage interacting pairs, and cluster galaxies with similar data.  We find that classifying the groups as dynamically ``old'' or ``young'', depending on whether or not at least one quarter of their members are early-type systems, is physical and consistent with past classifications of HCGs based on their atomic gas content. Dynamically ``old'' groups are more compact and display higher velocity dispersions than  ``young'' groups. Late-type galaxies in dynamically ``young'' groups have specific star formation rates (sSFRs), NUV-r, and mid-infrared colors which are similar to those of field and early stage interacting pair spirals. Late-type galaxies in dynamically ``old'' groups have redder NUV-r colors, as they have likely experienced several tidal encounters in the past building up their stellar mass, and  display lower sSFRs. We identify several late-type galaxies which have sSFRs and colors similar to those of elliptical galaxies, since they lost part of their gas due to numerous interactions with other group members. Also, 25\\% of the elliptical galaxies in these groups have bluer UV/optical colors than normal ellipticals in the field, probably due to star formation as they accreted gas from other galaxies of the group, or via merging of dwarf companions. Finally, our SED modeling suggests that in 13 groups, 10 of which are dynamically ``old\", there is diffuse cold dust in the intragroup medium. All this evidence point to an evolutionary scenario in which the effects of the group environment and the properties of the galaxy members are not instantaneous. Early on, the influence of close companions to group galaxies is similar to the one of galaxy pairs in the field. However, as the time progresses, the effects of tidal torques and minor merging, shape the morphology and star formation history of the group galaxies, leading to an increase of the fraction of early-type members and a rapid built up of the stellar mass in the remaining late-type galaxies.} ",
        "introduction": "It has become increasingly evident that interactions and  merging of galaxies have contributed substantially to their evolution, both in terms of their stellar population as well as their morphological appearance. Compact galaxy groups, with their high galaxy density and signs of tidal interactions among their members, are ideal systems to study the impact on environment to the evolution of galaxies. The Hickson Compact Groups (HCGs) are 100 systems of typically 4 or more galaxies in a compact configuration on the sky  \\citep{Hickson82}. They contain a total of 451 galaxies and  are mostly found in relatively isolated regions where no excess of surrounding galaxies can be seen, reflecting a strong local density enhancement. The HCGs occupy a unique position in the framework of galaxy evolution, bridging the range of galaxy environments, from field and loose groups to cores of rich galaxy clusters. The fact that the original selection of the HCG members did not include redshift information, led to the inclusion of interlopers among them, the most famous being NGC 7320 in  Stephan's Quintet (HCG  92).  This led to a debate as to whether compact groups  are line-of-sight alignments of galaxy pairs within loose groups, or filaments seen end-on \\citep{Mamon86, Hernquist95}. However, the detection of hot X-ray gas in $\\sim$75\\% of the HCGs by \\citet{Ponman96} implies that they reside in a massive dark matter halo and thus are indeed physically dense structures. Numerical simulations indicate that in the absence of velocity information, raising the minimum surface brightness criterion for the group used by Hickson would help eliminate  interlopers \\citep[see][]{McConnachie08}.  Because of the nature of these groups, the high density enhancements in addition to the low velocity dispersions ($\\sim$250 km s$^{-1}$), make them ideal to study the effects of galaxy interactions. \\citet{Hickson82} found that the majority of HCGs display an excess of elliptical galaxies, $\\sim$31\\% of all members compared to the field, while the fraction of spiral galaxies and irregular is only 43\\%, nearly a factor of two less of what is observed in the field. Optical imaging by \\citet{Mendes94} showed that  43\\% of all HCG galaxies display morphological features of interactions and mergers, such as bridges, tails and other distortions. Similar indications of interactions are seen in maps of the atomic hydrogen distribution in selected groups by \\citet{Verdes01}. Moreover, \\citet{Hickson89} found that the fractional distribution of the ratio of far-infrared (far-IR) to optical luminosity in HCG spiral galaxies is significantly larger than that of isolated galaxies, suggesting that for a given optical luminosity, spiral galaxies in groups have higher infrared luminosities. Comparison of HCG spirals with those in clusters of galaxies from \\citet{Bicay87} reveals that the distributions of the IR to optical luminosity, as well as the 60 to 100$\\mu$m far-IR color are similar. Finally, nuclear optical spectroscopy studies indicate that almost 40\\% of the galaxies within these groups display evidence of an active galactic nucleus (AGN, \\citealt{Martinez10,Shimada00}). All these clues are consistent with an evolutionary pattern where tidal encounters and the accretion of small companions by the group members, redistribute the gas content of the groups and affect the morphology of their members.  \\citet{Verdes01} and \\citet{Borthakur10} have proposed an evolutionary sequence for the HCGs based on the amount and spatial distribution of their neutral atomic gas. Using HI maps they classified the groups  into three phases based on the ratio of the gas content within the galaxies over the total observed in the group. According to their scenario, a loose galaxy group starts to contract under the gravity to form a more compact one. During this first phase the HI gas is still mostly found in the individual galaxies. Then as the group evolves, it enters the second phase and a fraction of the atomic hydrogen is extracted from the galaxy disks into the intragroup medium, probably due to tidal forces, and part of it becomes fully ionized. Finally, in the third phase, the dynamical friction leads to a decrease in the separation between the group members and the group becomes more compact. More than  80\\% of the HI originally in the disks of the galaxies has been displaced. Some of it is seen redistributed in a common envelope surrounding all groups members, and a fraction has likely been converted into molecular gas fueling star formation and/or accretion onto an AGN. A similar classification was proposed by \\citet{Johnson07} and \\citet{Tzanavaris10}. These authors separated the groups using the ratio of the M$_{\\rm HI}$ over the dynamical mass of the group, the so-called ``HI richness\", and found that it is correlated to the specific star formation (sSFR) of their galaxies. They also found that galaxies in gas-poor groups display colors representative of a normal stellar population and observed a bimodality in the mid-infrared colors, as well as in the sSFRs of their members. They suggested that this is caused by enhanced the star formation activity which lead the galaxies in groups to evolve faster. \\citet{Walker09} had concluded that a similar bimodality in the mid-IR is also observed in the colors of galaxies in the Coma Infall region.  A necessary step to determine the evolutionary state of HCGs, is the analysis of not just the morphology of the group members, but of their stellar population and star formation history. In our first paper (\\citealt{Bitsakis10}), we commenced exploring the properties of an initial sample of 14 HCGs in  mid-IR wavelengths and found that a large fraction of the early type galaxies in groups, have mid-IR colors and spectral energy distributions (SEDs) consistent with those of late-type systems. We suggested that this possibly stems from an enhanced star formation, as a result of accretion of  gas rich dwarf companions. We found no evidence that the star formation rate (SFR) and built up of stellar mass of late-type galaxies in groups is different from what is observed in early-stage interacting pairs, or spiral galaxies in the field. However, when we separated the groups according to the fraction of their early type members, they appear to differentiate from the control samples. This would  suggests an evolutionary separation of HCGs and can provide a better insight on the nature of these groups.  Despite the progress in the analysis of the properties of these groups, there are still several open questions. Is the bimodality of the mid-IR colors and sSFRs indeed linked with the evolutionary sequence of these galaxies, or it is also observed in other galaxies? How do the properties of HCG galaxies compare to those of galaxies in other environments? Is it really physical to classify the evolutionary stages of these groups also according to the fraction of their early-type systems? Is this classification meaningful in terms of galaxy properties and colors and does it agree with the other classifications? To answer  these questions we need  multi-wavelength data, as well as a theoretical model to fit their global SED, in order to obtain the best possible constraints on the physical parameters. In this paper we present the first such analysis for a large sample,  nearly one third of all Hickson groups, for  which we have retrieved observations from the  UV to the infrared part of the spectrum.  The structure of the paper is the following. We present our samples, the observations and the data reduction in Section 2. In Section 3, we describe the model used to fit the data as well as the basic physical parameters we can derive, along with their uncertainties. Our results are shown in Section 4, and our conclusions are summarized in Section 5. In an Appendix we provide additional information on10 early-type galaxies which have peculiar mid-IR colors and SEDs. ",
        "conclusions": "In this paper we have presented the first multi-wavelength analysis, from the UV to the infrared, of the 135 galaxies which are members of 32 Hickson Compact Groups and studied the influence of the group environment in their evolution. Based on the theoretical modelling of their SEDs we conclude the following:  \\begin{itemize}  \\item The classification of the evolutionary state of HCGs according to the fraction of their early-type members appears to be physical and is in general agreement with previous classifications. The study of their properties suggest that dynamically ``old'' groups are more compact and have higher velocity dispersions. They also display higher stellar masses than the ``young'' ones, while both have similar HI mass distributions. However, ``old'' groups have nearly an order of magnitude larger dynamical masses than ``young'' groups.  \\item The late-type galaxies in dynamically ``old'' groups display lower sSFRs since the multiple past interactions have already converted a fraction of their gas into stars increasing their stellar masses. This is also the main reason why these galaxies show redder NUV/optical colors than field spirals. However, there are few spiral galaxies in these groups which display even redder colors. They all have very small SFRs, similar to early-type systems. We speculate that tidal interactions must have stripped the gas out of their disk suppressing their star formation activity.  \\item Most early-type galaxies in  dynamically ``old'' groups, seem to migrate from the star-forming to the quiescent galaxy colors, even though a fraction of them ($\\sim$25\\%) display bluer colors and higher star formation activity than normal field ellipticals possibly due to gas accretion from other group members as well as merging of dwarf companions.  \\item Late-type galaxies in dynamically ``young'' groups have similar star formation properties to field spirals, as well as in early-stage interacting pairs.  \\item Even though nearly 46\\% of the HCG members have an optically identified AGN, we find no evidence of enhanced AGN activity at any stage of the group evolution, or the optical/mid-IR colors of the galaxies.  \\item Our analysis suggests that the reported lower density of galaxies in the IRAC color-color diagram is caused by the morphological natural bimodality of galaxies and it is similar to what is also observed in the UV-optical colors.  \\item Our SED model suggests that in 13 groups, 10 of which are dynamically ``old\", there is diffuse cold dust in the intragroup medium.  \\item There are 10 galaxies with mid-IR colors, SEDs, and star formation properties which are inconsistent with their optical classification as early-type systems. We suggest that 7 of them are likely nearly edge-on late-type galaxies which were misclassified due to their small size and dust extinction.    \\end{itemize}"
    },
    "1107/1107.1604.txt": {
        "abstract": "Motivated by the recent re-confirmation by CoGENT of the low-energy excess of events observed last year, and the recent improved limits from the XENON-100 experiment that are in contention with the CoGENT data, we re-examine the low mass neutralino region of the Minimal Supersymmetric Standard Model and of the Next-to-Minimal Supersymmetric Standard Model, both without assuming gaugino mass unification. We make several focused scans for each model, determining conservative constraints on input parameters. We then determine how these constraints are made increasingly stringent as we re-invoke our experimental constraints involving the dark matter relic abundance, collider constraints from LEP and the Tevatron, and then from flavour physics, as a series of successive $2\\sigma$ hard cuts.  We find that for both models, when all relevant constraints are applied in this fashion, we do not generate neutralino LSPs that possess a spin-independent scattering cross section in excess of $10^{-5}\\pb$ and a mass $7\\gev\\lsim \\mchi \\lsim 9\\gev$ that is necessary in order to explain the CoGENT observations. ",
        "introduction": "\\label{sec:intro}  Recent data from several direct detection (DD) dark matter (DM) experiments have significantly heightened interests in the search and identification of the elusive DM, that is predicted to dominate the matter component of the Universe, in the form of weakly-interacting massive particles (WIMPs). These include the observations made by the CoGENT experiment~\\cite{cogent}, a $p$-type point contact Germanium detector, whose enhanced sensitivity, due to an unprecedented combination of target mass and reduced electronic noise, revealed a possible excess of low energy events~\\cite{Aalseth:2010vx} that has recently been re-affirmed in June 2011~\\cite{Aalseth:2011wp}. If the exponential trend in these low recoil energy events is fully attributed to the elastic scattering of WIMP DM, then this signal is consistent with limits on the DM mass in the range $7\\gev\\lesssim\\mdm\\lesssim9\\gev$ and nucleon scattering cross section  in the range $4\\times10^{-41}\\,{\\rm cm}^{-2}\\lesssim\\sigman\\lesssim1.5\\times10^{-40}\\,{\\rm cm}^{-2}$, respectively~\\cite{Aalseth:2010vx, Aalseth:2011wp}. Remarkably, the CoGENT results favour a WIMP mass consistent with that indicated by the annual modulation count rate\\footnote{In their June 2011 paper, CoGENT also claim to observe evidence of an annual modulation in their data that is consistent with the DAMA/LIBRA results~\\cite{Aalseth:2011wp}. } measured by DAMA/LIBRA~\\cite{DAMA, Bernabei:2010mq}, assuming elastic scattering and that $\\sigman$ is spin-independent (SI)~\\cite{Petriello:2008jj, Chang:2008xa, Fairbairn:2008gz, Savage:2008er, Savage:2009mk}.  However, the majority of the region in $\\mdm$\\,-\\,$\\sigman$ parameter space favoured by CoGENT appears to be in contention with recent results from direct detection experiments, in particular, those recently reported by XENON-100~\\cite{XENON}, whose 2011 data consist of non-observations that appear to completely exclude the region favoured by CoGENT~\\cite{Aprile:2011hi}. The CoGENT and DAMA/LIBRA results are also in contention with the recent observation of two ``anomalous'' events by CDMS-II~\\cite{cdms, Ahmed:2009zw}, although the discrepancies here are much less extensive than with XENON-100, and such events are not statistically significant enough to be able to claim a detection of DM (see, e.g.,~\\cite{Fitzpatrick:2010em}). Some authors have previously concluded that the extent of the discrepancies between the results from CoGENT and DAMA and those from the other DD experiments may be partially reconciled by, for example, assuming a different proportion of elastic scattering events at DAMA/LIBRA that are channeled, adopting DM halo models (specifically pertaining to the escape velocity and velocity distribution of WIMPs), assuming a lower scintillation efficiency for Xenon, or assuming different couplings between WIMPs and either protons or neutrons (see e.g.,~\\cite{Fitzpatrick:2010em} and references therein).  Motivated by past experimental findings, several authors have investigated the parameter space of various particle physics models capable of generating light WIMP DM candidates that may potentially explain the CoGENT and DAMA/LIBRA observations (see, e.g.,~\\cite{Draper:2010ew, Das:2010ww, Vasquez:2010ru, Fornengo:2010mk, Cao:2011re}). However, many such investigations achieve this by conducting focused scans with finely-tuned parameters that result in points lying close to the CoGENT/DAMA favoured region, but provide a very poor fit to astrophysical and collider constraints.  In light of the recent release of the 2011 data from CoGENT and XENON-100, we follow a similar line and re-examine and directly compare the general behaviour of the neutralino low mass region of two popular particle physics models: the Minimal Supersymmetric Standard Model (MSSM) and the Next-to-Minimal Supersymmetric Standard Model (NMSSM).  We select to investigate the NMSSM in addition to the popular MSSM since in the NMSSM we expect that a lower mass lightest supersymmetric particle (LSP) is more easily allowed, relative to that say in the MSSM, owing to the weakening of the LEP bounds on the Higgs mass and couplings to Standard Model (SM) particles due to its extended Higgs sector.  To conduct our investigation we use the nested sampling (NS) algorithm, implemented in the Bayesian interface tool \\texttt{Multinest}~\\cite{multinest}, that is incorporated into the \\texttt{SuperBayeS}~\\cite{superbayes} and NMSPEC~\\cite{nmssmtool} packages, to perform scans focused on the low neutralino mass region when imposing experimental constraints in the likelihood function. However, unlike other recent studies, we then also impose the same constraints on the parameter space via a series of (2$\\sigma$) hard cuts on the resulting scan data. These hard cuts consist of an increasing number of constraints involving, firstly, those on the cosmological DM relic abundance, as inferred from cosmic microwave background (CMB) and galaxy cluster observations, secondly, from direct search limits inferred from particle colliders, e.g., LEP and the Tevatron, and, thirdly, from indirect constraints from flavour physics and the anomalous magnetic moment of the muon $\\gmtwo$.  For each set of hard cuts we highlight the parameter space for each of our models that is consistent with the limits on the LSP mass and SI nucleon scattering cross section that are imposed by DD limits.  The structure of this paper is as follows: In Sec.\\,\\ref{sec:procedure} we describe the details of the scanning procedure that we utilise to investigate the low mass sectors of our two particle physics models.  In Sec.\\,\\ref{sec:mssm} we present and discuss the results of our scan of the MSSM obtained when imposing our three sets of constraints from (i) the DM relic density, (ii) collider physics (excluding flavour physics) and (iii) indirect limits (flavour physics and $\\gmtwo$).  In Sec.\\,\\ref{sec:NMSSM} we present the results of our analogous scans on the NMSSM, and compare these results with those corresponding to the MSSM. Finally, in Sec.\\,\\ref{sec:summary} we summarise our conclusions.\\\\  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:summary}  In this paper we have re-examined the low mass neutralino region of both MSSM and NMSSM parameter space in light of the recent re-confirmation of a low energy excess of events observed by CoGENT, and the recent limits arising from the non-observations by XENON-100 that are in contention with the CoGENT data. To conduct our investigation we utilised focused NS scans when imposing a variety of experimental constraints, associated with (i) the cosmological DM relic abundance, (ii) direct collider search limits and (ii) collider physics relating to flavour physics and $\\deltagmtwo$, in the likelihood. Unlike the majority of previous studies we then re-invoked each of these constraints on the resulting scan data, this time, via a successive series of $2\\sigma$ hard cuts.  Firstly, regarding the MSSM, we relaxed the unification of gaugino masses in order to efficiently generate light, bino-like neutralino LSPs without violating collider constraints involving the invisible $Z$ width and the chargino mass.  We observed that for many configurations from our initial scan producing the neutralino LSP mass $\\mchi\\gsim5 \\gev$ the corresponding values of $\\abundchi$ were small enough to be within the WMAP $2\\sigma$ range. This could be achieved predominantly through efficient $t$-channel annihilations via stau exchange, which are enhanced for points possessing large $\\tanb$.  However, owing to LEP constraints on the slepton masses, when we invoked collider constraints as $2\\sigma$ hard cuts, only bino-like LSPs with $\\mchi\\gsim11 \\gev$ survived.  Moreover, when invoking the constraints from flavour physics, the minimum LSP mass was further increased to approximately $\\mchi\\gsim13 \\gev$, which is slightly heavier than the values within the favoured CoGENT region, and much lighter than the $28\\gev$ limit claimed by Vasquez {\\it et al.}~\\cite{Vasquez:2010ru}. We find that for these points it is extremely difficult to generate values of the SI elastic scattering cross section much larger than $10^{-5}\\pb$, owing primarily to the constraints on the $\\chi Hp$ coupling, which is enhanced for large $\\tanb$, coming from flavour physics, particularly the limit on $\\brbsmumu$.  We emphasise that this limit is particularly stringent within the framework of minimal flavour violation which we have adopted but can be fairly easily relaxed when one abandons it; see~\\cite{Foster:2005wb, Foster:2006ze}. However, we point out that solutions possessing such large values of $\\sigsip$, within our permitted mass range of $\\mchi\\gsim13 \\gev$, are also clearly excluded by the DD limits from XENON-100.  Regarding the NMSSM, once again, we conducted a focused NS scan in the low LSP mass region, where (non-flavour) collider constraints were invoked in the likelihood this time as $2\\sigma$ hard cuts. Once again, in order to efficiently generate light neutralino LSPs without violating LEP bounds on the invisible $Z$ width and the chargino mass, gaugino unification was relaxed. Whilst the majority of points surviving collider constraints were found to be bino-like, we found that a significant proportion of these points also evaded collider constraints by adopting a partial singlino component, made possible through the extended Higgs sector of the NMSSM. In such cases, we found that for light LSPs possessing $\\mchi\\lsim15 \\gev$ we were able to generate values of $\\abundchi$ small enough to lie within the WMAP $2\\sigma$ range through $s$-channel resonance annihilations via the lightest CP-even Higgs.  Consequently, in order to evade LEP bounds, in such cases the lightest Higgs had to be singlet-like, very much decoupled from the doublet Higgs. This could be achieved by adopting the limit $\\lambda\\rightarrow0$, in which case bino-dominated LSPs are favoured, or by finely-tuning $A_\\lambda$, generally resulting in a combination of bino-singlino LSP configurations and a hierarchy between $\\lambda$ and $\\kappa$ where $\\kappa\\ll\\lambda$ for $\\lambda\\gg0$ that we explicitly described, or by finely-tuning $A_\\kappa$. However, for $\\mchi\\gsim15 \\gev$, we found that points can generate $\\abundchi$ small enough to survive our cuts by virtue of additional $t$-channel annihilations via very light sleptons, possible through our relaxation of the universality of slepton masses, or in some cases through $s$-channel annihilations via the pseudoscalar Higgs.  We found that the points from our initial scan generated values of $\\sigsip$ up to approximately $2\\times10^{-5}\\pb$. However, we found that many of these points possessed too large a bino fraction to survive the flavour physics constraints, particularly that coming from $\\brbsmumu$, that we later re-invoked as $2\\sigma$ hard cuts, finding results very similar to those obtained in~\\cite{Das:2010ww} where $\\sigsip \\lsim 10^{-6}\\pb$. We also highlight that many of these points with $\\mchi\\gsim10\\gev$ are excluded by DD limits from XENON-100, like many of the points discussed by other authors, including~\\cite{Vasquez:2010ru,Cao:2011re}, that possess $\\sigsip$ as large as $10^{-2}\\pb$, but are generated without invoking experimental constraints as $2\\sigma$ hard cuts, as we do, in order to omit points that grossly violate individual constraints.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1107/1107.5302_arXiv.txt": {
        "abstract": "The return to the quiescent state of the Anomalous X-ray pulsar \\xte\\ following its 2003 outburst represents a unique opportunity to probe the surface emission properties of a magnetar.  The quiescent emission of \\xte\\ is composed of two thermal components, one arising from the whole star surface, and the other from a small warm spot on it. By modeling the magnitude and shape of the pulse profile in narrow spectral bands, we have been able to constrain the physical characteristics and geometrical parameters of the system: the two angles that the line of sight and the spin axis make with respect to the warm spot axis ($\\psi$ and $\\xi$ respectively), the angular size of the spot, and the overall surface temperature distribution.  Our modeling accounts for the general relativistic effects of gravitational redshift and light bending near the stellar surface, and allows for local anisotropic emission.  We found that the surface temperature distribution on the neutron star is consistent with the expectations of a dipole magnetic field configuration; the local radiation requires a pencil-beamed emission pattern, suggesting the presence of a magnetized atmosphere.  For a typical value of the radius, R=13~km, the viewing parameters (symmetric for an interchange between $\\psi$ and $\\xi$), range from $\\psi=\\xi=38^{\\circ}$ to ($\\psi,\\xi$)=($52^{\\circ}$,$29^{\\circ}$). These angles are consistent with those obtained by modeling the AXP in outburst, with uncertainty contours reduced by a factor of 2.5. ",
        "introduction": "\\xte\\ is an isolated neutron star (NS) belonging to the class of the Anomalous X-ray Pulsars (AXPs); these objects, together with the Soft Gamma-ray repeaters (SGRs), are believed to be magnetars: isolated neutron stars whose thermal emission and occasional outbursts are powered by their extremely strong magnetic fields (Duncan \\& Thompson 1992; Thompson \\& Duncan 1995).  Originally one of the thousands of the faint X-ray sources cataloged by ROSAT, the variable nature of \\xte\\, was revealed by the outburst of 2003, with its sudden increase in X-ray luminosity by a factor $\\sim 100$, decaying on a time-scale of years.  The source rotation period and its derivative were consequently measured and found to be $P=5.54\\,\\rm s$ and $\\dot{P}=1.1-2.1\\times10^{-11}\\,\\rm s/s$. These timing properties imply a magnetic field $B_{\\rm dip}\\sim 3\\times10^{14}\\,\\rm G$, confirming the magnetar classification of the source (Ibrahim et al. 2004, Gotthelf at. al. 2004). The source was monitored repeatedly for more than 7 years with several X-ray observatories, up to the return to quiescence (Gotthelf \\& Halpern 2005, 2007; Halpern \\& Gotthelf 2005;  Bernardini et al. 2009;  Albano et al 2010).  These studies presented a unique opportunity to probe the emission mechanisms of a strongly magnetized NS by taking advantage of the flux evolution during its decay. While analysis of phase-averaged spectra alone cannot uniquely distinguish among competing emission models, the addition of the steady change of the spectrum and pulse profile over time greatly increases the diagnostic power.  Perna \\& Gotthelf (2008) developed a detailed emission model for the energy-dependent pulse profile of \\xte\\, following its outburst. This model, which was tailored to the specific surface emission distribution in the post-outburst phase, can take on any viewing geometry, includes the general relativistic effects of light deflection and gravitational redshift, and allows for local anisotropic emission. The application of this model to the first 4 sets of \\XMM\\ data acquired during the temporal evolution of the flux from \\xte\\ following the outburst (September 2003 - September 2004), provided a constraint on the underlying emission geometry and radiation properties of this transient magnetar in its post-outburst phase.  In this paper we present the results of the modeling of the spectral and timing data of \\xte\\ obtained with 3 combined \\XMM\\ pointings (September 2009) upon the return of the source to quiescence (see Figure \\ref{fig:fluxevol}, last point). Indeed, our goal here is that of studying the properties of the quiescent emission of this magnetar, which carries information on the surface temperature distribution of the star, and hence on its magnetic topology. We model these data using a modified version of the emission model by Perna \\& Gotthelf (2008), updated to include: {\\em (a)} the presence of quiescent emission from the full surface of the star, {\\em (b)} the changed and reduced emission from the region heated by the outburst.  The new data and its spectral and timing analysis is reported in \\S2. In \\S3, we discuss the properties of the quiescent emission. The theoretical emission model is described in detail in \\S4, and the results of its application to the data are given in \\S5, following with a discussion in \\S6.   \\begin{figure} \\centering \\includegraphics[scale=0.45,angle=270]{evflux_sum.ps} \\caption{{\\em XMM} (red) and {\\em Chandra} (blue) 0.5--10 keV flux measurements of \\xte\\ following its outburst.  It is evident from the data that \\xte\\ has now reached quiescence. The dashed line represents the X-ray flux level ($\\sim7.5\\times10^{-13}\\,\\ergscm$) as recorded by $ROSAT$, Einstein, and $ASCA$ before the outburst onset (Ibrahim et al. 2004, Gotthelf et al. 2004). The analysis presented in the paper is performed over three close \\XMM\\ pointings represented by the last data point.} \\label{fig:fluxevol} \\end{figure} ",
        "conclusions": "The decline of the post-outburst emission of \\xte\\  has allowed us to perform a detailed spectral and timing analysis of the emission of this source in its quiescent state. The main results from the modeling described herein can be summarized as follows:  \\begin{enumerate}  \\item The temperature distribution on the surface of the NS is consistent with the expectations of a dipole magnetic field configuration.  We remark that our analysis does not guarantee that a dipolar magnetic field  represents the only possible solution; in principle, other magnetic field configurations and beaming patterns can be ad-hoc produced and tested for this source. However, a dipolar field is the simplest magnetic field configuration for a pulsar, and one that yields a single peak, nearly sinusoidal (as observed in \\xte), for a wide range of viewing angles. Hence, although not formally unique from a mathematical point of view, the proposed solution is physically motivated, and the fact that it provides such a good match to the data yields confidence that it is indeed a reasonably good representation of the quiescent emission of this transient magnetar. \\item The pulsed fraction of the NS emission requires an anisotropic, pencil-type radiation pattern, which is a likely indication of the presence of a magnetized atmosphere on the NS surface.  For a { $\\cos^{n}\\delta$} emission profile, we performed a timing analysis for $n$ in the range 0-2.5 (in steps of 0.5), and radii between 9-15~km (in steps of 1~km).  We found that no match to the PF data could be obtained for $n=0$ (isotropic radiation pattern), no matter the value of the radius, while for $n=2.5$, a good match to the PFs could be found even for $R=9$~km (smaller values of the radius are increasingly allowed as $n$ increases).  For each value of $n$ in between there is a range of radii which are not allowed by the PF data (e.g. for $n=1$ no good match could be found for $R\\le 12$~km). Therefore, our analysis has clearly demonstrated how a detailed {\\em a priori} knowledge of the spectral distribution and emission pattern as a function of ($\\theta, \\phi$) on the entire surface of the star has the power to allow a radius constraint for the NS.  \\item The overall emission geometry is constrained by identifying likelihood regions in the $\\psi-\\xi$ parameter space. The most significant range according to our fits is consistent with the results of the earlier, post-outburst analysis of Perna \\& Gotthelf (2008), but the $3\\sigma$\u03c3 allowed region is now reduced by a factor of 2.5. We note that the hot-spot axis used in the earlier work coincides with the dipole axis in the current study.  A spectral and timing analysis of the post-outburst emission of \\xte\\, has been performed also by Albano et al. (2010). They fitted the first 7 observations using a three-temperature model. The X-ray emission is produced in a globally twisted magnetosphere, and parameterized by the twist angle, the electron velocity and the seed photon temperature. In the the last two sets, the flux is approaching the quiescent level, and only two temperatures are required to fit the X-ray spectrum.  For these observations, they found that the angle between the magnetic axis and the rotation axis is $\\xi= {30.0^\\circ}^{+12.3}_{-20.0}$ in the 6th observation, and $\\xi= {22.7^\\circ}^{+16.4}_{-20.0}$ in the 7th, 6 months later, while the angle between the line of sight and the rotation axis was found to be $\\psi= {153.9^\\circ}^{+19.6}_{-16.0}$ in the 6th and $\\psi= {145.8^\\circ}^{+16.4}_{-9.5}$ in the 7th observation.  While the values of the angle $\\xi$ are roughly within the range of what we find, the angle $\\psi$ is inconsistent with ours. A direct comparison between our modeling and theirs is however not possible due to some fundamental differences in the basic assumptions. Our quiescent emission is assumed to be thermal, and we leave the anisotropy level of the local radiation as a free parameter. On the other hand, the twisted magnetosphere model has a well determined angular radiation pattern. The NS surface temperature is assumed to follow a dipolar-like pattern in our model, while in theirs it is assumed to be constant.  Furthermore, our modeling includes the general relativistic effect of light deflection (which heavily influences the pulsed fractions and hence the viewing geometry that we infer), while their modeling does not.  \\end{enumerate}  Our study has allowed us to establish some important properties of the quiescent emission of \\xte\\,. The spectral analysis has shown evidence for the presence of a very small hot spot, only a couple of degrees in size. This emission component dominates the X-ray flux above about 1 keV. Interestingly, a small hot region is also characteristic of the quiescent emission of another (transient) AXP, \\cxou\\, (Israel et al. 2007; Skinner et al. 2006). The most natural interpretation, in the case of \\xte\\, as well as in the case of \\cxou\\,, is the association of this warmer region to a concentration of magnetic field lines, since the surface heat flux of a magnetar scales sensitively as $B^{4.4}$ (Thompson \\& Duncan 1996). It is then not surprising that this region coincides with the location of the outburst, as our analysis for \\xte\\, has shown.  Our constraints on the viewing geometry could in principle be combined with those obtained from radio measurements following the outburst, to further restrict the allowed parameter space. Camilo et al. (2007a) showed that the radio and X-ray locations were aligned, within some rather large uncertainty due to the small number of counts of their {\\em Chandra} data. Camilo et al. (2007b) used radio polarimetry to fit for the angle angle between the magnetic field and the rotation axis ($\\xi$ in our notation), and for the angle between the magnetic field axis and the line of sight ($\\alpha_{min}=\\psi-\\xi$ in our notation), finding two possible configurations: one with $\\xi=70^\\circ$ and $\\alpha_{min}\\sim 20^\\circ-25^\\circ$, and another with $\\xi=4^\\circ$ and $\\alpha_{min}=4^\\circ$.  Polarization observations of \\xte\\, were also used by Kramer et al. (2007) to constrain the viewing geometry of this pulsar.  They identified a main pulse and a mid pulse, produced in different locations on the star. Their best fitted angles (with the same notation as above) were $\\xi=44^\\circ$ and $\\alpha_{min}=39^\\circ$ for the main pulse, and $\\xi=76^\\circ$ and $\\alpha_{min}=6^\\circ$ for the inter pulse.  Given some discrepancies between the different angle estimates from radio data, and the lack of a firm and accurate association between the radio and X-ray emitting regions, we refrain in the current work from using the radio angle determinations to restrict our allowed parameters range. However, we take the opportunity to emphasize the importance of simultaneous radio-X ray observations should another outburst occur (either in this or in another magnetar).  Our analysis has allowed us to explore the consistency of the data with the expectations for the temperature distribution on the NS surface if the $B$ field is dominated by a dipolar component.  The angles $\\xi$ and $\\psi$ that provide a good match to the observed pulsed fractions (Fig.~3 and Table~2) also ensure that the profile remains single-peaked in the lowest energy band, dominated by the surface/dipolar component. Combinations of the beaming parameter and the NS radius can be found for which our model provides an excellent match to spectra, pulsed fractions and, more generally, to the full pulse profile.  Within the context of magnetars, being able to infer the magnetic topology is important for a number of reasons. Firstly, measurements of the magnetic field strength from $P$ and $\\dot{P}$ make the implicit assumption of a dipolar field; substantial departures from this configuration will result in biased estimates. Second, the magnetic field topology in magnetars also plays an important role in the resulting outbursting behaviour (Thompson \\& Duncan 1995, 1996), in that it regulates the intensity, frequency and location on the NS surface of the starquakes, and hence of the outbursts (Perna \\& Pons 2011).  Ultimately, we note how the combined spectral/timing analysis that we performed has the potential to constrain the compactness ratio of the neutron star (or equivalently the radius, for a given mass). However, obtaining this type of constraint requires an a priori knowledge of the spectral and radiation pattern of the local emission.  This is particularly difficult in the case of highly magnetized objects, since the local emission depends on the local strength and orientation of the $B$ field, which in turn would need to be determined as a part of the fitting procedure. This analysis,  which has begun for NSs with magnetic fields in the $10^{12}-10^{13}$~G range (Ho \\& Mori 2008), will become possible for magnetars once atmospheric models for arbitrary field orientations  (Lloyd 2003a,b; Ho et al. 2007) are extended to strengths $B\\sim 10^{14}-10^{15}$~G.  \\begin{figure*} \\begin{center} \\begin{tabular}{cc} \\includegraphics[angle=-90,scale=.43]{d2.ps}  \\\\ \\includegraphics[angle=-90,scale=.41]{fig7a.ps} \\\\ \\end{tabular} \\caption{\\textit{Upper panel}: \\XMM\\ phase-averaged spectrum of \\xte\\ obtained with the model presented in the text, for the specific case of R=13 km, D=3.3 kpc, and the best fit angles $\\psi^{*}=38^{\\circ}$ and $\\xi^{*}=38^{\\circ}$. Model residuals are shown in the lower panel. The parameters for the best fit model are reported in Table \\ref{tab:sp_faceon}. The data show the summed spectra from the three observations. \\textit{Lower panel}: Same as the upper panel, but for unfolded source spectra. The contribution of the different model component is also shown: surface (black dotted line), warm spot (red dotted line).} \\label{fig:spec13} \\end{center} \\end{figure*}"
    },
    "1107/1107.2008_arXiv.txt": {
        "abstract": "We revisit previously developed analytic models for defect evolution and adapt them appropriately for the study of semilocal string networks. We thus confirm the expectation (based on numerical simulations) that linear scaling evolution is the attractor solution for a broad range of model parameters. We discuss in detail the evolution of individual semilocal segments, focusing on the phenomenology of segment growth, and also provide a preliminary comparison with existing numerical simulations. ",
        "introduction": "Cosmic strings are line-like concentrations of energy that can arise as topological defects in cosmological phase transitions in the early universe \\cite{Kibble80,VilShell}. Their notable popularity in the 80's was due to the belief that they could be responsible for structure formation, an alternative scenario to inflationary perturbations \\cite{Hawking82,Mukhanov82,Linde82}.  At the end of the 90's, however, it was realised \\cite{AlbrBat,AlbrBat1} that strings cannot be solely responsible for the observed structure, which was further confirmed by the direct observations of acoustic peaks in the CMB, in accordance with the inflationary prediction. This led to a shift of attention away from cosmic strings.  However, recent developments in fundamental theory have led to a resurgence of interest in the subject \\cite{KibbleRev,PolchIntro}.  First, it was realised that the production of cosmic strings is much more generic than previously thought in a wide class of cosmological models based on supergravity \\cite{Jeannerot}.  In addition, new exciting developments in string phenomenology \\cite{PolchBranes,KiritsisRev,QuevedoRev} allowed the study of cosmological inflation from brane constructions, and this led to a generic picture of hybrid inflation ending in a phase transition that produces cosmic (super)strings \\cite{BMNQRZ,MajumDavis,DvalTye}.  These objects are much more complex than ordinary field theory strings and their properties depend sensitively on fundamental high-energy physics parameters of the underlying theory \\cite{JoStoTye,PolchProb,PolchStab}. Thus, not only it was realised that string-like defects are much more generic objects that could play a subdominant role in structure formation, but, also, their potential observability could open a window into physics at the highest energy scales.  Over the past few years, there has been significant progress in understanding the properties, cosmological evolution and possible observational signatures of cosmic superstrings (for a recent review see \\cite{CopPogVach}).  However, there remain certain types of well-motivated string objects for which their physics, and most notably their cosmological evolution, remains relatively unexplored. For example, there is yet no detailed self-consistent model describing the evolution of cosmic strings carrying currents, even though such strings have been predicted in situations where more than one cosmological phase transition is considered \\cite{Brandenberger:1996zp}, and they also arise naturally in SUSY models that predict strings  \\cite{Davis:1997bs,Achucarro:2002jg}. String currents also appear in the effective 4D description of higher dimensional strings, as the string position in the internal dimensions is described by worldsheet scalar fields giving rise to currents.  Another example is semilocal strings \\cite{VachAchuc}, arising in SUSY GUT inflationary models \\cite{Urrestilla:2004eh} and the corresponding D3/D7 brane inflation models \\cite{D3D7semiloc}. These models are a natural extension of usual inflationary models, in which the only extra ingredient needed is doubling a hypermultiplet. This simple change renders the properties of the strings very different; as we will discuss in the following section, these semilocal strings are non-topological, they have different stability properties ~\\cite{Vachaspati:1991dz,Hindmarsh:1991jq} and even their CMB constraints are less stringent than those for usual cosmic strings~\\cite{Urrestilla:2007sf}.  More recently, monopole networks and hybrid networks (where finite string segments have monopoles attached to their ends) have also been studied \\cite{MartAch,Hybrid}. Both in the gauge and in the global case, the hybrid networks tend to annihilate very quickly, usually in less than a Hubble time after formation, although they may survive for a longer period in particular circumstances. There is an outstanding hybrid-like case, where local strings attached to global monopoles. This, at least naively, is the starting point for the much more interesting semilocal case which we discuss here.  String evolution is a difficult problem involving physics from different energy  scales, which is often not fully understood. There have been many different approaches for obtaining the properties of string networks, trying to negotiate between computational power and accuracy, between numerical and analytic approaches. Numerical simulations approaches include field theoretic simulations, Nambu-Goto approximations, and the so-called unconnected segment toy models \\cite{AlbrBat}. All these approaches have their virtues and their limitations, and progress has been possible through combination  of results from different fronts. Another complementary approach is that of analytic velocity-dependent one-scale (VOS) models \\cite{VOS96,VOS02}. These abandon a `statistical physics' description of the network in favour of a `theromodynamics' one, characterising the network by a small number of physically meaningful macroscopic quantities such as correlation lengths and root-mean-squared (RMS) velocities, whose evolution equations can be calculated from the microscopic equations. This approach has the obvious advantages of simplicity, tractability and applicability beyond the dynamical ranges accessible to simulations, but it comes at a cost---the averaging process (through which one deduces the macroscopic evolution equations from the microscopic ones) requires the introduction of a (relatively small) number of phenomenological parameters, whose values can only be inferred by using numerical simulations for calibration.  In what follows we discuss such an analytic model for the evolution of semilocal string networks. We study the behaviour of the network as a whole, starting from the premise that it can be treated as a network of local strings attached to global monopoles. We also describe the evolution of individual semilocal segments, discussing under what conditions these segments can grow---a process that has been clearly identified in numerical simulations. Finally, we provide a preliminary comparison between the analytic model and existing numerical simulations, focusing on the overall network properties and the dependence on the coupling parameter. A more detailed comparison, discussing in detail the distribution of the semilocal segments and its time evolution, is left for a follow-up paper. ",
        "conclusions": "We have extended previously developed analytic models for defect evolution, applying them to the study of semilocal string networks. We have shown that a linear scaling evolution (analogous to the well-known one for cosmic strings) is the attractor solution for a broad range of model parameters. Our results therefore provide supporting evidence for those of existing numerical simulations, where confirmation of the presence of this solution is restricted to the limited dynamical range available.  We have discussed in some detail the evolution of individual semilocal string segments, focusing on the interesting phenomenology of segment growth and considering two different possible scenarios thereof. As expected (and as seen in simulations) we find that small segments tend to shrink and decay while large ones can grow and merge with others. Which (if any) of the two scenarios we have considered is the valid one is an issue that can only be resolved with future and more detailed simulations.  Finally we have used the results of existing simulations in order to obtain a very preliminary comparison with our analytic results. This endeavour is quite limited by the fact that the available simulations have not measured the defect velocities, which play a key role in VOS-type models. We found some encouraging results, but it is quite clear that a more detailed comparison, using expanding-universe simulations that explore a wider range of the space of model parameters (notably the coupling parameter $\\beta$), will be needed in order to provide a meaningful calibration of the model. We shall address these issues in a subsequent publication."
    },
    "1107/1107.0657_arXiv.txt": {
        "abstract": "{We study the bright Seyfert 1 galaxy Mrk~509 with the Reflection Grating Spectrometers (RGS) of XMM-Newton using the RGS multi-pointing mode of XMM-Newton for the   first time in order to constrain the properties of the outflow in this object.} {We want to obtain the most accurate spectral properties from the 600~ks spectrum of Mrk~509 which has excellent statistical quality.} {We derive an accurate relative calibration for the effective area of the RGS, derive an accurate absolute wavelength calibration, improve the method for adding time-dependent spectra and enhance the efficiency of the spectral fitting by two orders of magnitude.} {We show the major improvement of the spectral data quality due to the use of the new RGS multi-pointing mode of XMM-Newton. We illustrate the gain in accuracy by showing that with the improved wavelength calibration the two velocity troughs observed in UV spectra are resolved.} {} ",
        "introduction": "Outflows from Active Galactic Nuclei (AGN) play an important role in the evolution of the super-massive black holes (SMBH) at the centres of the AGN, as well as on the evolution of the host galaxies and their surroundings. In order to better understand the role of photo-ionised outflows, the geometry must be determined. In particular, our goal is to determine the distance of the photo-ionised gas to the SMBH, which currently has large uncertainties. For this reason we have started a large monitoring campaign on one of the brightest AGN with an outflow, the Seyfert 1 galaxy \\object{Mrk~509} \\citep{kaastra2011}. The main goal of this campaign is to track the response of the photo-ionised gas to the temporal variations of the ionising X-ray and UV continuum. The response time immediately yields the recombination time scale and hence the density of the gas. Combining this with the ionisation parameter of the gas, the distance of the outflow to the central SMBH can be determined.  The first step in this process is to accurately determine the physical state of the outflow: what is the distribution of gas as a function of ionisation parameter, how many velocity components are present, how large is the turbulent line broadening, etc. Our campaign on Mrk~509 is centred around ten observations with XMM-Newton spanning seven weeks in the Fall of 2009. The properties of the outflow are derived from the high-resolution X-ray spectra taken with the Reflection Grating Spectrometers \\citep[RGS,][]{denherder2001} of XMM-Newton. The time-averaged RGS spectrum is one of the best spectra ever taken by this instrument, and the statistical quality of this spectrum  can be used to improve the current accuracy of the calibration and analysis tools. This gives rise to challenges for the analysis of the data. The methods developed here also apply to other time-variable sources. Therefore we describe them in some detail in this paper.  For this work we have derived a list of the strongest absorption lines in the X-ray spectrum of Mrk~509, and we perform some simple line diagnostics on a few of the most prominent features. The full time-averaged spectrum will be presented elsewhere \\citep{detmers2011}. ",
        "conclusions": "\\subsection{Wavelength scale and spectral resolution}  By using five different methods, we have obtained a consistent wavelength scale with an accuracy of about 1.8~m\\AA, four times better than the nominal wavelength accuracy of RGS. The only slightly deviant diagnostic remains the neutral gas in our galaxy (\\ion{O}{i} and \\ion{N}{i}), which are off from the mean by $8.8\\pm 3.8$~m\\AA\\ (2.3$\\sigma$). The systematics due to blending (for \\ion{O}{i}) and the lack of sufficiently reliable RGS data for \\ion{N}{i} may be slightly larger than we have estimated.  Our work shows that in cases like ours a dedicated effort to improve the wavelength scale is rewarding. As we showed, at this level of accuracy, Doppler shifts due to the orbit of Earth around the Sun become measurable. It also allows us to use X-ray lines as tools to study the dynamics of the outflow and the Galactic foreground at a level beyond the nominal resolution of the instrument (see Sect.~\\ref{sect:agn}).  \\subsection{Effective area corrections}  The corrections that we have derived for the relative effective area of both RGS detectors, as well as for the relative effective area of both spectral orders are at the few percent level but should be taken into account for the high statistical quality spectrum that we consider here. This is due to the fact that there are data gaps for each individual RGS, but each time at a different wavelength range. Without such a correction, there would be artificial jumps at the boundaries of these regions of the order of a few percent. This leads any spectral fitting package to attempt to compensate for this by enhancing or decreasing astrophysical absorption lines or edges, in the hunt for the lowest $\\chi^2$ value.  A disadvantage of our method is of course that the ``true'' effective area and hence absolute flux is known only to within a few percent, and the deviations are wavelength-dependent. In our subsequent spectral fitting \\citep{detmers2011} we compensate for this by using a spline for the AGN continuum. This is clearly justified, as one might worry about the intrinsic applicability of e.g. pure power-law or blackbody components when the flux can be measured down to the percent level accuracy. The ``true'' continua may contain all kinds of features including e.g. relativistically broadened lines \\citep[e.g.][]{branduardi2001,sako2003} that can be discerned only in the highest-accuracy data, such as the present one.  \\subsection{Combination of spectra}  The reduction of the combined response matrices of our 40 spectra with 2 Gb of memory into a single spectrum and response file of only 8 Mb, a reduction by more than two orders of magnitude, allows us to make complex spectral fits. In the most sophisticated models used by \\citet{detmers2011} the number of free parameters approaches 100, and without this efficient response matrix the analysis would have been impossible.  \\subsection{Final remarks}  Due to the high statistical quality of our data, the binning problems mentioned in Sect.~\\ref{sect:statistics} became apparent. In many RGS spectra that have been obtained during the lifetime of XMM-Newton, the statistical quality of the spectra is not good enough to detect this effect. Still, this could imply that when users obtain a fit with reduced chi-squared of order unity, the fit is formally not acceptable because an ideal fit would yield a smaller reduced chi-squared. Clearly, here more work needs to be done to alleviate the problem, but due to the discretisation of wavelength in the detector corresponding to the natural size of the CCD pixels there is some natural limitation to what can be done. Even in the best observations there is always some small scatter in the satellite pointing, making some rebinning necessary.  Overall, however, the calibration of RGS is at a fairly advanced stage. Only for exceptionally good spectra like the Mrk~509 spectrum that we discussed here the sophisticated tools developed in this paper are needed, although also for weaker sources our improvements can be beneficial for the analysis. Our tools are publicly available as auxiliary programs ({\\sl rgsfluxcombine}, {\\sl rgs\\_fmat}) in the SPEX distribution."
    },
    "1107/1107.0182_arXiv.txt": {
        "abstract": "{We report on very long baseline interferometry (VLBI) observations of the radio emission from the inner region of the Crab Nebula, made at 1.6\\,GHz and 5\\,GHz after a recent high-energy flare in this object.  The 5 GHz data have provided only upper limits of 0.4\\,milli-Jansky (mJy) on the flux density of the pulsar and 0.4\\,mJy/beam on the brightness of the putative flaring region. The 1.6 GHz data have enabled imaging the inner regions of the nebula on scales of up to $\\approx 40\\arcs$. The emission from the inner ``wisps'' is detected for the first time with VLBI observations.  A likely radio counterpart (designated ``C1'') of the putative flaring region observed with Chandra and HST is detected in the radio image, with an estimated flux density of $0.5\\pm 0.3$\\,mJy and a size of 0{\\farcs}2--0{\\farcs}6. Another compact feature (``C2'') is also detected in the VLBI image closer to the pulsar, with an estimated flux density of $0.4\\pm0.2$\\,mJy and a size smaller than 0{\\farcs}2. Combined with the broad-band SED of the flare, the radio properties of C1 yield a lower limit of $\\approx 0.5$\\,mG for the magnetic field and a total minimum energy of $1.2 \\times 10^{41}$\\,ergs vested in the flare (corresponding to using about 0.2\\% of the pulsar spin-down power). The 1.6 GHz observations provide upper limits for the brightness (0.2\\,mJy/beam) and total flux density (0.4\\,mJy) of the optical Knot 1 located at 0{\\farcs}6 from the pulsar. The absolute position of the Crab pulsar is determined, and an estimate of the pulsar proper motion ($\\mu_\\alpha = -13.0 \\pm 0.2$\\,mas/yr, $\\mu_\\delta = +2.9 \\pm 0.1$\\,mas/yr) is obtained. ",
        "introduction": "The inner region of the Crab Nebula contains a complex and dynamic structure with a number of elliptical ripples (\"wisps\") varying in flux on timescales of days and moving outward with speeds of up to $0.7\\,c$ \\citep{tanvir1997,hester2002} in the optical, while radio measurements with the VLA yield somewhat lower speeds of $\\approx 0.3\\,c$ \\citep{bietenholz2004}. The origin of the wisps is not yet known, but they are generally thought to be associated with instabilities downstream of the shock in the wind from the pulsar that powers the nebula \\citep{gallant1994}.   The wind derives its power from the pulsar spin-down injecting an energy of $\\sim 5 \\times 10^{38}$\\,erg/s into relativistic particles and magnetic field \\citep{rees1974,kennel1984}. The particles are thought to be mostly electrons and positrons, accelerated to relativistic energies, with possibly a lesser number of ions \\citep{arons1994}.  The exact mechanism by which the spin-down energy is released, transported, and dissipated remains poorly understood \\citep[cf.,][]{chedia1997,hester1998,begelman1999,lyutikov2003,spitkovsky2004,komissarov2004} -- with shocks, plasma instabilities, pair acceleration, Poynting flux, and magnetic field all suggested to affect the evolution of the wind.  Fast variability possibly related to abrupt energy releases in the nebula was first reported in the radio on time scales of $\\sim 1$ day, corresponding to an emitting region of $\\le 0{\\farcs}1$ in size \\citep{matveyenko1975,matveyenko1979}. On 22 September 2010, a gamma-ray flare was discovered for the first time from the direction of the Crab Nebula with the AGILE pair-production telescope \\citep{tavani2011}, offering a rare opportunity to follow in detail the evolution of a flaring region in the nebula. The flare was confirmed with Fermi/LAT \\citep{abdo2011} indicating a variability time-scale as short as 12~hours \\citep{balbo2011}.  The flare was not detected in the keV regime with Swift/XRT \\citep{evangelista2010} and INTEGRAL \\citep{ferrigno2010a}. The Fermi/LAT monitoring had further shown an abrupt decrease of the gamma-ray flux on 23 September, with the emission returning to its pre-flare level. Based on the short duration of the high-energy flare, it has been suggested \\citep{komissarov2010} that the flaring material is located in the optical Knot 1 separated by $\\approx 0{\\farcs}6$ from the pulsar \\citep{hester1995}. However, subsequent high-resolution imaging with Chandra (on 28 September, \\citeauthor{tennant2010} \\citeyear{tennant2010}) revealed three stationary compact ($\\approx 1^{\\prime\\prime}$ in size) regions (labelled A,B, and C) of enhanced emission located on the X-ray bright ring. Similar structures are are also visible in an HST image of the Crab Nebula taken on 02 October \\citep{caraveo2010}.  These knots present themselves as a viable site for the flaring emission \\citep{tavani2011}, and their compactness may explain naturally the non-detections in observations with lower spatial resolution where the extended nebula emission dominates.  The Chandra and HST observations of the knots indicate a strong, rapid, and efficient energy release in the flaring material and pose the question of the relation between the optical, X-ray, and gamma-ray emission.  Given the position of the knots within the nebula (located not too far off the jet axis and slightly closer than the innermost wisp), it is not clear whether the flaring region is associated with the jet or with the equatorial outflow. Deciding on these issues relies on detailed imaging of the emitting region, which is best achieved in the radio regime.  We present here high-resolution wide-field radio images of the flaring knot and the central region of the nebula, obtained with very long baseline interferometry (VLBI) observations at a wavelength of 18\\,cm. These observations probe a wide range of angular scales, from $\\sim 10$ milliarcseconds to $\\sim 40$ arcseconds, detect the flaring region, and reveal, for the first time, the intricate morphological structure of the radio emission within about 20 parsecs of the pulsar. The observation and data reduction are described in Section~2. The resulting images are presented in Section~3, and discussed in the context of physics and evolution of the flare in the Crab Nebula.   \\begin{table} \\caption{Antenna parameters} \\label{tb:antennas} \\begin{center} \\begin{tabular}{l|cccr}\\hline\\hline \\multicolumn{1}{c|}{Antenna} & \\multicolumn{1}{c}{$D$} & \\multicolumn{1}{c}{$\\lambda_\\mathrm{obs}$} & SEFD & \\multicolumn{1}{c}{$B_\\mathrm{min}$--$B_\\mathrm{max}$} \\\\ & \\multicolumn{1}{c}{[m]} & \\multicolumn{1}{c}{[cm]} & \\multicolumn{1}{c}{[Jy]} & \\multicolumn{1}{c}{[km]} \\\\ \\hline Effelsberg      & 100     & 6,\\,18 &  20,19 &  266--~~853   \\\\ Westerbork      &  66$^a$ & 6,\\,18 &  60,30   &  266--1002  \\\\ Jodrell Bank    &  25     & 6,\\,18 &  320,320 &   17--1388  \\\\ Medicina        &  32     & 6,\\,18 &  170,600 &  757--1429  \\\\ Onsala          &  25     & 6,\\,18 &  600,320 &  601--1429  \\\\ Torun           &  32     & 6,\\,18 &  220,230 &  637--1441  \\\\ Yebes           &  40     & 6      &  90  & 1313--2152  \\\\ Cambridge$^b$   &  32     & 6      &  136 &  197--1313  \\\\ Darnhall$^b$    &  25     & 18     &  356 &   17--1404  \\\\ Knockin$^b$     &  25     & 18     &  356 &   67--1441  \\\\ Hartebeesthoeck$^c$ &  26 & 18 &  450 & 7453--8525  \\\\\\hline \\end{tabular} \\end{center} \\small {Notes:}~$a$ -- for Westerbork operating as a phased array; $b$ -- MERLIN antennas; $c$ -- Hartebeesthoeck was added ad hoc in the observation at 18\\,cm, while participating in the e-VLBI session on November 5. Column designation: $D$ -- antenna diameter; $\\lambda_\\mathrm{obs}$ -- observing wavelength used; SEFD -- system effective flux density describing the antenna sensitivity at corresponding wavelength; $B_\\mathrm{min}$--$B_\\mathrm{max}$ -- range of the telescope separations from a given antenna (for European telescopes, baselines to the South African antenna in Hartebeesthoeck are not considered). \\normalsize \\end{table}  \\begin{table} \\caption{Target and calibrator sources} \\label{tb:targets} \\begin{center} \\begin{tabular}{clc}\\hline\\hline \\multicolumn{2}{c}{Object/Coordinates} & $\\sigma_\\mathrm{pos}$[mas] \\\\ \\hline \\multicolumn{2}{c}{{ Crab-A}$^a$}  &  \\\\ $\\alpha_\\mathrm{J2000}$ & ~~05{\\hh}34{\\mm}32{\\farcsh}332 & $\\sim$10 \\\\ $\\delta_\\mathrm{J2000}$ & +22{\\dg}00{\\arcm}52{\\farcs}45 & $\\sim$10 \\\\ \\multicolumn{2}{c}{{J0518+2054}$^b$} & \\\\ $\\alpha_\\mathrm{J2000}$ & ~~05{\\hh}18{\\mm}03{\\farcsh}824510 & 0.12\\\\ $\\delta_\\mathrm{J2000}$ & +20{\\dg}54{\\arcm}52{\\farcs}49739 & 0.12 \\\\\\hline \\end{tabular} \\end{center} \\small {Notes:}~$\\sigma_\\mathrm{pos}$ -- position errors in right ascension and declination. References: $a$ -- \\citeauthor{tavani2011} \\citeyear{tavani2011} and this work; $b$ -- \\citeauthor{petrov2008} \\citeyear{petrov2008}. \\normalsize \\end{table} ",
        "conclusions": "The VLBI observations of the Crab nebula have yielded accurate positional and photometric information about the pulsar itself and the flaring region, as well as marginal morphological and photometric information about the inner wisps of the nebula.  \\subsection{Astrometry and proper motion of the pulsar}  The position of the pulsar is measured relative to the position of the phase-reference calibrator, J0518+2054. The position of J0518+2054 has been determined with an accuracy of 0.12 mas, both in right ascension and declination \\citep{petrov2008}. The calibrator image obtained from the VLBI data shows no significant structure that could affect the calibrator phases and introduce additional positional errors. The total flux density recovered in the calibrator image is $290.1 \\pm 1.8$\\,mJy, and the peak flux density is $271.0 \\pm 1.0$\\,mJy/beam, further underlying the compactness of the structure. The calibrator image has a convolving beam of $18 \\times 12$\\,mas and an r.m.s. noise of 0.8\\,mJy/beam (and the respective peak-to-noise SNR of 340), thus limiting the angular size of the calibrator to $\\le 1.1$\\,mas and introducing an error of $0.06$\\,mas to the positional measurements \\citep{lobanov2005}.  For the calibrator, we obtain the following J2000 position (the errors listed are formal errors of the fit): \\[ \\alpha_\\mathrm{J2000} = \\textrm{05:18:03.8245014} \\pm 0.0000014\\,, \\] \\[ \\delta_\\mathrm{J2000} = \\textrm{20:54:52.497662} \\pm 0.000028\\,. \\] Analysis of the high-resolution image of the pulsar yields an apparent J2000 position (the errors are formal errors of the fit): \\[ \\alpha_\\mathrm{J2000} = \\textrm{05:34:31.9383014} \\pm 0.0000081\\,, \\] \\[ \\delta_\\mathrm{J2000} = \\textrm{22:00:52.17577} \\pm 0.00018\\,. \\] Taking into account the positional errors of the calibrator listed in Table~\\ref{tb:targets}, more conservative estimates of the errors of the pulsar position are given by $\\Delta_\\alpha = 0.18$\\,mas (0{\\farcsh}000012) and $\\Delta_\\delta = 0.22$\\,mas. These estimates do not include errors due tropospheric and ionospheric delays and errors from the calibrator-target phase transfer normally considered in measurements made with dedicated VLBI astrometry observations \\citep[cf.,][]{guirado1995,martividal2008}. Based on the analysis of these factors \\citep{guirado1995}, we can expect that their contribution to the positional errors should be within $\\approx 0.05\\, \\mathrm{FWHM\\,[mas]}\\,\\Delta\\theta\\mathrm{[deg]}$ milliarcseconds, with $\\Delta\\theta$ giving the calibrator-target distance. For the FWHM of the full resolution image, one can therefore expect that the positional errors can reach $\\Delta\\alpha = 3.4$\\,mas and $\\Delta\\delta = 0.45$\\,mas (reflecting a much better north-south resolution of the image).  \\begin{figure*}[t!] \\includegraphics[height=1.0\\textwidth,angle=-90,clip=true]{al-fig07.eps} \\caption{Image of the radio emission (green contours; multi-scale CLEAN image, convolved with a 0{\\farcs}7 beam; contours start at $\\mathrm{SNR} = 2$ and increase with an interval of 0.5) overlaid with images of the optical (left panel, 02 October 2010; \\citeauthor{caraveo2010} \\citeyear{caraveo2010}) and X-ray (right panel, 28 October 2010; \\citeauthor{ferrigno2010b} \\citeyear{ferrigno2010b}) emission in the flaring region. The radio emission from the flaring region C1 is offset from the knot A \\citep{tavani2011} identified in the optical and X-ray emission. The overall appearance of the optical and radio emission is similar in the flaring, with a clear ``knee-like'' structure seen in the optical and trailed by a similar structure in the radio. It should also be noted that the apparent discrepancy in appearance of C1 and C2 in this figure and in the image in Fig.~3 (convolved with a $0{\\farcs}15$ beam) results from the fact that the component C2 is compact ($\\le 0{\\farcs}2$ in size) while the size of C1 can be as large as $0{\\farcs}6$ (see discussion in Sect.~\\ref{sc:flreg})} \\label{fg:evn-hst} \\end{figure*}   The position measurement made from the VLBI data can be combined with the Crab pulsar positions obtained from pulsar timing \\citep{taylor1993,han1999}, yielding the following estimate of the pulsar proper motion:  \\[ \\mu_\\alpha = -13.0 \\pm 0.2\\quad \\textrm{mas/yr}\\,, \\] \\[ \\mu_\\delta = +2.9 \\pm 0.1\\quad \\textrm{mas/yr}\\,. \\] This estimate agrees well with recent estimates made from a long range of HST observations \\citep{ng2006,kaplan2008}.    \\subsection{Flaring region} \\label{sc:flreg}  The flaring components C1 and C2 are both very weak, detected at $\\textrm{SNR}\\le 4$, which makes it difficult to provide robust estimates of their properties. Table~\\ref{tb:flare} lists our best estimates of the positions, flux densities, and sizes of the two components, based on Gaussian fits to the single-scale and multi-scale image (with the latter providing a somewhat better account of the extent of the emitting regions). We use these positions to estimate brightness temperature of the emission in respective regions. The resulting values indicate brightness temperature in excess of $\\approx 4000$\\,K (it should be noted here that estimating such low values of brightness temperature is enabled by the presence of short baselines provided by the MERLIN antennas). The upper limit on the size of C1 corresponds to the maximum size of a 0.5\\,mJy feature that could be detected in the data.  \\begin{table}[ht] \\caption{Properties of compact flaring components} \\label{tb:flare} \\begin{center} \\begin{tabular}{l|c|c}\\hline\\hline & C1 & C2 \\\\\\hline R.A. [s]      & $32.2271 \\pm 0.0011$ & $32.03851\\pm 0.00052$ \\\\ Dec. [\\arcs]  & $~53.096  \\pm 0.014$ & $~52.8600 \\pm 0.0070$ \\\\ $S_\\mathrm{1.6\\,GHz}$ [mJy] & $0.5 \\pm 0.3$ & $0.4 \\pm 0.2$ \\\\ $\\theta$ [\\arcs] & 0{\\farcs}2 -- 0{\\farcs}6 & $\\le 0{\\farcs}2$ \\\\ $T_\\mathrm{b}$ [K] & 1000 -- 5000 & $\\ge 4000$  \\\\ \\hline \\end{tabular} \\end{center} Notes: $S_\\mathrm{1.GHz}$ -- component flux density; $\\theta$ -- estimated size (C2 is unresolved, and a limit on the size is estimated from the SNR following \\cite{lobanov2005}; for C1, the upper limit on the size is given by the detection limit in the data); $T_\\mathrm{b}$ -- brightness temperature calculated for the respective size and flux estimates (note that short baselines of MERLIN ($10^{-2}$--$10^{-1}\\mathrm{M}\\lambda$) provide sensitivity to emission with brightness temperature greater than $\\approx 1000$\\,K). \\end{table}  An overlay of the radio image of the flaring region with an HST image from 2 October \\citep{caraveo2010} and a Chandra image from 28 October \\citep{ferrigno2010b} is shown in Fig.~\\ref{fg:evn-hst}. The radio image is convolved with a 0{\\farcs}7 beam to further emphasise weak extended emission. This overlay shows that the radio emission traces well the inner wisp and ``sprite'' (region B), commonly identified in the optical/X-ray images. The radio component C1 is located close to a ``knee-like'' region (region A) detected in the optical and X-ray images and believed to be associated with the flare \\citep{tavani2011}. {This morphology may also be reflected by the radio emission, but establishing this relation clearly requires better quality of radio imaging on arcsecond scales.}   \\begin{figure*}[t!] \\includegraphics[height=0.95\\textwidth,angle=-90]{al-fig08.eps} \\caption{\\label{fig:SED} The spectral energy distribution of the nebula and the flaring component C1.  The data on the nebula have been compiled in \\cite{meyer2010}. Details on the modelling of the synchrotron nebula are provided in that reference as well. In addition to the nebula, the SED of the feature C1 is shown together with a time-dependent model of the flare (solid red line), with spectra calculated at the epoch of the flare maximum ($t_\\mathrm{f}$) and one day before ($t_\\mathrm{f-1d}$) and after ($t_\\mathrm{f+1d}$) the maximum } \\end{figure*}  \\subsection{Origin of the radio knots}  A significant offset of both C1 and C2 from the jet axis is evident. If these two features are associated with the jet (and are not part of the wisps produced in the equatorial outflow), it may be speculated that they are located on a ``wall'' formed by interaction of the rotating outflow with the ambient material in the nebula.  The offsets $\\Delta\\,r$ of C1 and C2 with respect to the jet axis, $z$, directed at a position angle $\\psi_\\mathrm{jet} \\approx 130\\dg$ (estimated from the image in Fig.~\\ref{fg:flare}) can be represented by a power-law expansion $\\Delta\\,r \\propto z^m$, which yields $m \\approx 0.7$. A jet viewing angle of $60\\dg$ \\citep{cheng2000,komissarov2010} is assumed in these calculations. The expansion derived is consistent with the one expected for the jet boundary in a magnetically confined flow \\citep{fendt1997,fendt2001} with a weakly evolving poloidal magnetic field component. One can assume that the jet boundary is defined by the outermost flux surface \\citep[cf.,][]{fendt2001} and that the collimation of the flow begins near the light cylinder, $R_\\mathrm{lc}$, with an initial radius of the flow of $\\approx 10\\,R_\\mathrm{lc}$. It would then require an electric current $I_\\mathrm{p} \\approx 4 \\times 10^{14}\\,\\mathrm{A}$, a magnetic field $B \\approx 4\\times 10^{12}$\\,G, and a magnetic flux $\\Psi_\\mathrm{p} \\approx 6 \\times 10^{24}\\,\\mathrm{G}\\,\\mathrm{cm}^2$ (all agreeing well with earlier estimates made for the Crab pulsar; cf., \\citeauthor{hankins2007} \\citeyear{hankins2007}; \\citeauthor{hirotani2006} \\citeyear{hirotani2006}) to explain the locations of the knots C1 and C2, assuming that these knots are located near the jet boundary. We note that the value of $I_\\mathrm{p}$ is larger than the Goldreich-Julian current, which can be explained by pair cascades yielding particle densities well in excess of the Goldreich-Julian density \\citep{arendt2002}. With the jet expansion derived, the jet should have local half-opening angles of $\\phi_\\mathrm{C1} = 27\\dg$ and $\\phi_\\mathrm{C2} = 37\\dg$ at the locations of the radio knots. For the HST Knot 1, the respective half-opening angle is $\\phi_\\mathrm{HST\\,1} = 56\\dg$, which is close to the value of $60\\dg$ expected for this region \\citep{komissarov2010}. These arguments lend further support to the suggestion that the radio knots are connected to the outflow.  However, providing a firm conclusion on this matter would certainly require a closer followup of another prominent flare in the Crab nebula.    \\subsection{Connection to the high-energy flare}  Based on the observational results summarised above, the faint and compact radio features C1 and C2 could in principle be related to the gamma-ray flare, even if they are not co-spatial with the active regions observed in the optical and X-ray bands. In the following, we consider a scenario in which the brighter feature C1 is related to the gamma-ray flare and discuss the consequences.  The spectral energy distribution (SED) of C1 is shown in Fig.~\\ref{fig:SED} together with the nebula emission \\citep[see][for details of the data]{meyer2010}. This SED combines the flux density measured at 1.6\\,GHz, the upper limit obtained at 5\\,GHz, and the X-ray flux of the knot A treated here as a (soft) upper limit on the X-ray emission from the immediate vicinity of C1. Assuming that the underlying particle distribution present in the region C1 has the same spectral slope ($\\mathrm{d} n/\\mathrm{d}\\gamma = n_0 \\gamma^{-1.60}$) as the electrons responsible for radio emission in the nebula, the resulting spectrum matches quite well both the radio as well as the gamma-ray emission across 14 orders of magnitude in energy.   The observed variability time-scale of about 24~hrs \\citep{balbo2011} constrains the size of the emitting region to be smaller than $r < t_\\mathrm{var} c= 3\\times 10^{15}\\,\\mathrm{cm}\\, (t_\\mathrm{var}/10^5\\,\\mathrm{sec})= 0{\\farcs}1\\,(t_\\mathrm{var}/10^5\\,\\mathrm{sec})$. This is consistent with the observed extension of C1.  Assuming that the gyro-radius $r_\\mathrm{g} < ct_\\mathrm{var}$ and that gamma-rays are produced via synchrotron emission, a lower limit on the B-field is set by $B\\ge 0.5\\,\\mathrm{mG}\\,(E/100\\,\\mathrm{MeV})^{1/3}(t_\\mathrm{var}/100\\,\\mathrm{ksec})^{-2/3}$. A conservative upper limit on the magnetic field of $\\approx 100$\\,mG follows from the argument that the total energy in the magnetic field should not exceed the spin down power of the pulsar integrated over the duration of the flare.  The equipartition magnetic field (assuming minimum total energy) is $B_\\mathrm{eq}=3\\,\\mathrm{mG}\\,(t_\\mathrm{var}/100\\,\\mathrm{ksec})^{-6/7} (E/100\\,\\mathrm{MeV})^{1/7}$. The corresponding total minimum energy ($w_\\mathrm{tot} = w_\\mathrm{e} + w_\\mathrm{B} = 1.2\\times 10^{41}$\\,ergs).  The injection power required to generate the burst is $\\dot w\\approx w_\\mathrm{tot}/t_\\mathrm{var}=1.2\\times 10^{36}\\,\\mathrm{ergs/s} = 0.2$\\,\\% of the available spin-down power. Even though this is a comfortable margin, deviations from the minimum energy requirement are limited to roughly one order of magnitude.  Given the limited knowledge about the dynamics of the feature C1, it is difficult to estimate directly any effects related to relativistic beaming. A kinematic estimate of the Doppler factor $\\mathcal{D}=\\gamma^{-1}(1-\\beta \\cos \\vartheta)$ by adopting a jet viewing angle $\\vartheta = 30^\\circ$ favoured in several recent works \\citep[{\\em e.g.},][]{ng2004,komissarov2004,komissarov2010} and assuming that the jet material moves at a speed similar to the range of speeds observed in the wisps. This yields Doppler factors in a 1.4--1.8 range.  An upper limit on the $\\mathcal{D}$ can also be derived based upon the non-observation of an accompanying flare at very high energy gamma-rays. The contemporaneous observations using the VERITAS and MAGIC \\citep{mariotti2010,ong2010} air Cherenkov telescopes rule out any substantial deviations of the flux in the energy regime between 100 GeV and TeV. Taking the relative flux errors of roughly 10~\\% as an upper limit on any additional flaring component, we can estimate an upper limit on $\\mathcal{D}<\\mathcal{O}(100)$. The actual upper limit depends critically on the geometry. For the sake being conservative, we assume a synchrotron self-Compton type scenario neglecting additional photon fields.  These considerations (spectrum and size) {are consistent with the hypothesis that the region C1 detected in the radio image is connected with the high-energy flare in the Crab nebula. The unusually long duration of higher flux state during this event may receive further support from the detection of C2, which could be a second plasma condensation ejected from the pulsar after the main flare in September 2010 and leading to emission enhancement in the high-energy regime.  The observed positional offset between the features C1 and A and the lack of obvious optical/X-ray counterparts for the feature C2 raise general questions of the relation between the emission detected in all three bands and the nature of the flaring activity in the Crab nebula. In the absence of better quality data, these questions remain open. As a matter of speculation, one can suggest that the flaring activity may be occurring in a relatively thin layer close to the jet boundary, thus highlighting the jet edges where the path length through the emitting material is larger (indeed, the locations of the features A and C can also be reconciled with the same jet expansion as derived above for C1 and C2).  Such an activity may originate, for instance, from the magnetic reconnection or interaction of the jet with the ambient medium.  However, in the absence of more detailed observational data, it is clearly not feasible to provide a reliable judgement on this matter.  In this respect, we also would like to emphasise that further improvement of the quality of radio imaging on arcsecond scales (which should be expected after the e-MERLIN becomes fully operational) would provide much better reliability for imaging of such extended and complicated objects as the central regions of the Crab nebula. Therefore, engaging in a detailed monitoring program of the Crab nebula with the EVN and e-MERLIN combined would be a highly desirable option for a followup of any future strong flaring event in this object."
    },
    "1107/1107.5558_arXiv.txt": {
        "abstract": "The plasma $\\beta$ (the ratio of the plasma pressure to the magnetic pressure) of a system can have a large effect on its dynamics as high $\\beta$ enhances the effects of pressure anisotropies. We investigate the effects of $\\beta$ in a system of stacked current sheets that break up into magnetic islands due to magnetic reconnection, which is analogous to the compressed heliospheric current sheet in the heliosheath. We find significant differences between systems with low and high initial values of $\\beta$. At low $\\beta$ growing magnetic islands are modestly elongated and become round as contraction releases magnetic stress and reduces magnetic energy. At high $\\beta$ the increase of the parallel pressure in contracting islands causes saturation of modestly elongated islands as island cores approach the marginal firehose condition. Only highly elongated islands reach finite size. The anisotropy within these islands prevents full contraction, leading to a final state of highly elongated islands in which further reconnection is suppressed. The elongation of islands at finite $\\beta$ is further enhanced by reducing the electron-to-ion mass ratio, to more realistic values. The results are directly relevant to reconnection in the sectored region of the heliosheath where there is evidence that elongated islands are present, and possibly to other high $\\beta$ systems such as astrophysical accretion flows and the magnetosphere of Saturn. ",
        "introduction": "At the outer edges of the solar system, the solar wind pushes up against the interstellar medium. Since the solar wind is moving at supersonic speeds, it forms a shock known as the termination shock. The region between the termination shock and the interstellar medium is referred to as the heliosheath. Flapping of the heliospheric current sheet produces sectored magnetic fields \\citep{Smith01} that have a significant latitudinal extent. There has been research suggesting that the current sheets between the sectored fields found in the heliosheath are compressed to the point that collisionless reconnection begins to occur, resulting in the formation of magnetic islands \\citep{Drake10,Czechowski10}. A turbulent magnetohydrodynamic (MHD) model of the reconnection of the sectored fields has also been proposed \\citep{Lazarian09} although we will argue later that the \\emph{Voyager} data are inconsistent with this hypothesis.  An important question, however, is whether the conventional treatment of collisionless reconnection \\citep{Shay07} is valid in the heliosheath, where it was suggested that the pick-up ion (PUI) population increases the plasma pressure compared with values at 1 AU \\citep{Zank99, Richardson08,Wu09}. Although both \\emph{Voyager} spacecraft are currently taking data in the heliosheath, the energy range of the detectors does not cover the PUIs, so it is difficult to make a reliable estimate of the value for $\\beta$, the ratio of the plasma pressure to the magnetic pressure \\citep{Richardson08}. Global MHD simulations suggest, however, that $\\beta$ varies from $8$ to $0.5$ between the termination shock and the interstellar medium with the highest $\\beta$ just downstream of the termination shock \\citep{Drake10}. Although this simulation does not include a separate pick-up ion population, it provides a rough estimate for the expected values for $\\beta$ and motivates the range of $\\beta$ in our study. In this study we investigate the impact of $\\beta$ on the dynamics of reconnection and the formation of magnetic islands relevant to the sectored heliosheath.  To begin reconnection and island formation a current sheet needs to be compressed to approximately the ion inertial scale $d_i = c/\\omega_{\\text{pi}}$, where $c$ is the speed of light and $\\omega_{\\text{pi}}$ is the ion plasma frequency, \\citep{Cassak05,Yamada07}. At this point the current sheet becomes unstable to the collisionless tearing mode. Upstream of the termination shock the heliospheric current sheet has a thickness of around $10,000\\text{ km}$ \\citep{Smith01} and is predicted to compress to around $2500\\text{ km}$ just downstream of the shock. In the upstream region the plasma density measured at \\emph{Voyager 2} is around $0.001\\text{cm}^{-3}$, corresponding to an ion skin depth of around $7,200\\text{ km}$.  In the downstream region the density is compressed to around $0.003\\text{cm}^{-3}$, corresponding to an ion skin depth of around $4200\\text{ km}$. Thus, the compression of the current sheets across the termination shock should trigger collisionless reconnection in the heliosheath and some \\emph{Voyager 1} and \\emph{2} observations support this hypothesis \\citep{Opher11}.  In our system we will be examining symmetric current sheets with no guide field (initial out-of-plane magnetic field). In an analytic study \\cite{Brittnacher95} showed that when $\\rho_i/w_0 \\approx 1$, the fastest growing linear mode occurs at $kw_0 \\approx 0.5$, where $k$ is the wavenumber of the tearing mode, $w_0$ is the half width of the current sheet, and $\\rho_i$ is the ion gyroradius.  Since $\\rho_i=\\sqrt{\\beta_i}d_i$, where $\\beta_i$ is the plasma beta based on the ion pressure, the current sheet thickness is comparable to that in our simulation.  Due to the collisionless nature of the plasma, pressure anisotropies ($P_\\parallel \\neq P_\\perp$, where $\\parallel$ and $\\perp$ are determined with respect to the magnetic field) can form during reconnection. Fermi acceleration in contracting islands \\citep{Drake06a} and adiabatic cooling as $B$ decreases due to conservation of the magnetic moment $\\mu \\propto v_\\perp^2/B \\propto P_\\perp/B$ both drive $P_\\parallel > P_\\perp$. When an anisotropy is formed with $P_\\parallel > P_\\perp$ the tension in bent magnetic fields weakens. The fluid momentum equation with an anisotropy becomes \\begin{equation} \\rho \\frac{d\\mathbf{v}}{dt} = -\\mathbf{\\nabla}\\left(P_\\perp + \\frac{1}{8\\pi} B^2 \\right) + \\mathbf{\\nabla}\\cdot\\left[\\left(1 - \\frac{\\beta_\\parallel-\\beta_\\perp}{2}\\right) \\frac{\\mathbf{B}\\mathbf{B}}{4 \\pi} \\right] \\text{,} \\end{equation} where $\\rho$ is the mass density of the plasma, $\\mathbf{v}$ is the bulk velocity, and $\\mathbf{B}$ is the magnetic field. For $\\beta_\\parallel = \\beta_\\perp$ the pressure equation reduces to the standard MHD equation. When $\\beta_\\parallel > \\beta_\\perp$ the tension force is reduced. At high $\\beta$ this reduction of tension force is noticeable for even slight anisotropies in the pressure. For $\\beta_\\parallel-\\beta_\\perp$ large enough, the tension force drops to zero, or even becomes negative. Since the tension of field lines acts as a restoring force for Alfv\\'en waves in standard MHD, the negative sign causes this oscillation to become an instability known as the firehose instability \\citep{Parker58} for: \\begin{equation} \\label{firehose} \\beta_\\parallel - \\beta_\\perp > 2 \\text{.} \\end{equation} This instability is fueled by the free energy contained in the pressure anisotropy. The firehose instability causes magnetic field lines to kink, which eventually relieves the pressure anisotropy by causing scattering.  Alternatively, when $\\beta_\\parallel-\\beta_\\perp$ is negative and large enough in magnitude, other instabilities can occur. The mirror-mode instability and the ion cyclotron instability both occur when $P_\\perp > P_\\parallel$. For larger $\\beta_\\parallel$ the mirror-mode becomes unstable at smaller values of $|\\beta_\\parallel-\\beta_\\perp|$ than the ion cyclotron mode, so the marginal mirror-mode criterion acts as the boundary between the stable and unstable regions. Based on fluid theory assuming $T_e = T_i$ \\citep{Hasegawa69}, the mirror-mode instability occurs when \\begin{equation} \\label{mirrormode} \\beta_\\perp - \\beta_\\parallel > \\frac{\\beta_\\parallel}{\\beta_\\perp} \\text{.} \\end{equation}  There are also kinetic modifications that can be made to the marginal instability criteria for firehose, mirror-mode, and ion cyclotron which make them more accurate. Although a rigorous analytic theory is not available, there are models that approximate the instability very well \\citep{Hellinger06, Bale09}. However, for simplicity we will just consider the conditions based on fluid theory.  In this study, we simulate several stacked current sheets similar to the compressed sectored heliospheric fields and associated current sheets, and follow the development of reconnection and islands.  We implement this system in a two-dimensional particle-in-cell (PIC) code and vary the temperature of the background plasma to test the dependence on $\\beta$. We observe that in finite $\\beta_e$ systems ($\\beta_e > 0.5$), very elongated islands form as opposed to the modest-aspect-ratio islands found at low $\\beta_e$ ($\\beta_e < 0.5$), where $\\beta_e$ is the $\\beta$ based on the electron pressure. At high $\\beta$, the increased $P_\\parallel$ due to the Fermi reflection of electrons within islands saturates the normal modest-aspect-ratio islands. Fermi reflection in highly elongated islands is less efficient because of the increased bounce time of the electrons so these islands are able to reach finite amplitude. At late time, however, even these elongated islands exhibit anisotropy instabilities, from Fermi reflection of both ions and electrons. As a result, late-time magnetic islands remain highly elongated and do not become round as in the low-$\\beta$ regime. This result has significant implications for the structure of islands that would be measured in the heliosheath. Although $\\beta_e$ is rather moderate in the heliosheath, we find a mass ratio dependence suggesting long islands for a broad range of $\\beta_e$ in realistic mass ratios. A large $\\beta$ however may be necessary to sustain the elongation of these islands. ",
        "conclusions": "The magnetic islands that reach a significant amplitude are much more elongated at high $\\beta_e$ than at low $\\beta_e$. These elongated islands should be found even for moderate values of $\\beta_e$ at realistic mass ratios. Island elongation is caused by the suppression of the shorter wavelength tearing modes by pressure anisotropies ($P_\\parallel > P_\\perp$) that develop due to the Fermi acceleration of electrons. Later in time the plasma develops pressure anisotropies of both ions and electrons that are limited by the firehose and Weibel instabilities. A Weibel mode develops that kinks the magnetic field lines. In the regime with a real mass ratio we would expect even longer islands to form, where multiple wavelengths of the firehose instability could develop. At late time the fraction of points unstable to the firehose instability saturates, and the anisotropy is confined between the mirror-mode and firehose instability boundaries. The long islands persist due to the low requirement of anisotropy to reach the marginal firehose condition at high $\\beta$. For even small anisotropies the tension in the magnetic fields is removed.  When encountering magnetic islands in the heliosheath, we predict the formation of similar extended, sausage-shaped islands rather than the more round islands found in low-$\\beta$ simulations \\citep{Drake10}. The cores of these islands should also be at the marginal firehose condition, so the magnetic tension that drives them to become round vanishes. We would thus expect these sausage shapes to persist long after the islands have ceased growing, and thus could be found even in regions where reconnection is no longer occurring.  \\begin{figure} \\noindent\\includegraphics[width=3.0in]{lambdacompare.eps} \\caption{\\label{lambdacompare} Distribution of $\\lambda$ at $t= 110\\Omega_{\\text{ci}}^{-1}$ for the (a) $\\beta=0.2$ case and (b) $\\beta=4.8$ case. The dotted lines are at $\\lambda = 90^\\circ$ and $270^\\circ$ where we expect to find peaks in the distribution.} \\end{figure}  These elongated islands exhibit signatures that can be seen in \\emph{Voyager} data. In particular, \\emph{Voyager} measures all three components of the magnetic field. Of particular interest for the explorations of islands that grow in the ecliptic plane is the angle $\\lambda =\\tan^{-1}\\left(B_T/B_R\\right)$, where $B_T$ and $B_R$ are the azimuthal and radial magnetic fields, respectively. $\\lambda = 90^\\circ$ and $270^\\circ$ correspond to the azimuthal unreconnected sectored heliosheath magnetic fields. Deviation of $\\lambda$ from $90^\\circ$ and $270^\\circ$ indicates some process is distorting the sectored field. \\emph{Voyager} data show the distribution of $\\lambda$ is peaked in the two azimuthal directions, $\\lambda = 90^\\circ$ and $270^\\circ$ \\citep{Opher11}. These peaks are significantly broader in the heliosheath than upstream, indicating that reconnection or another mechanism is disturbing the heliosheath field. The observed \\emph{Voyager} distribution of $\\lambda$ is consistent with that found in high-$\\beta$ simulations\\citep{Opher11}. Since the islands are elongated, the magnetic fields tend to remain primarily in the azimuthal direction even well after the islands begin to interact with each other. Round islands, such as would be expected from an MHD model or a low $\\beta$ kinetic model, are not consistent with observations since they produce much broader $\\lambda$ distributions. Thus, MHD reconnection \\citep{Lazarian09} in the heliosheath seems to be ruled out. Shown in Figure\\,\\ref{lambdacompare} is the distribution of $\\lambda$ from the simulations at $\\beta=0.2$ (Figure\\,\\ref{jzbeta}(a)), and $\\beta=4.8$ (Figure\\,\\ref{jzbeta}(c)) at $t=110 \\Omega_{\\text{ci}}^{-1}$. The high $\\beta$ simulation which has elongated islands retains the two peaks at $\\lambda=90^\\circ$ and $\\lambda=270^\\circ$. The long islands have a larger magnetic field in the azimuthal direction than the radial, resulting in peaks in the $\\lambda$ distribution, but the shorter islands become round having a magnetic field with similar strength in both directions, resulting in a broad distribution in $\\lambda$. The loss of tension in a finite $\\beta$ plasma prevents the complete release of magnetic energy that would be expected in an MHD model. A complete understanding of the $\\beta$ dependence of magnetic islands is essential in order to obtain reliable signatures that can be compared with \\emph{Voyager} data.  In this work there was no out-of-plane guide magnetic field. In the heliospheric current sheet, the magnetic field rotates from one direction to the other keeping a constant magnitude rather than passing through zero \\citep{Smith01}. A guide field would cause the center of the islands to have a much lower $\\beta$ since the magnetic field does not go to zero. Because of this magnetic field, we would not expect the Weibel instability to develop. In real systems there is frequently a guide field, so this would be worth further investigation.  In astrophysical accretion disks, reconnection plays a role in determining the saturation of the magnetorotational instability (MRI) \\citep{Sano04}. The saturation of MRI is strongly dependent on the dissipation of the magnetic field due to reconnection. For the high $\\beta$ in accretion disks, suppression of the most strongly growing small islands may significantly impact the saturation of the MRI. Since $\\beta$ is typically larger than $100$ in these structures, the only surviving islands would be so long that it is likely that much of the magnetic free energy would not be dissipated. Further, since the MRI requires magnetic tension, the absence of tension could limit the development of the instability.  \\cite{Sharma06} perform a simulation showing an enhancement of the growth of MRI due to anisotropies with $P_\\perp > P_\\parallel$ , which enhances the magnetic tension, caused by $\\mu$ conservation as a magnetic field develops. They do not capture the physics of reconnection and Fermi acceleration in magnetic islands that would generate anisotropies with $P_\\perp < P_\\parallel$, which removes magnetic tension. These two competing sources of anisotropy, both affect the tension and thus the growth of the MRI. The relative importance of these mechanisms needs to be explored.  Reconnection at high $\\beta$, although relatively rare in the terrestrial magnetosphere, is also found in the magnetosphere of Saturn \\citep{Masters11}. Magnetic islands were discovered in a region where $\\beta$ is larger than 10. The $\\beta$ dependence of the growth of finite-sized magnetic islands may lead to a better understanding of these findings.  Since the development of elongated islands requires only moderate $\\beta$, we expect to see the development of longer islands than expected in lower $\\beta$ systems such as the magnetosphere. In contrast to \\cite{Karimabadi05}, these longer islands can grow to a large enough size to play a role in magnetospheric dynamics. Since the $\\beta$ of the magnetosphere is not exceptionally large it is unlikely that the persisting anisotropy is enough to keep the islands from eventually becoming round.  \\smallskip  The computations were performed at the National Energy Research Scientific Computing Center.  1,1           Top"
    },
    "1107/1107.5072_arXiv.txt": {
        "abstract": "We use three-dimensional hydrodynamical simulations to study the rapid infall phase of the common envelope interaction of a red giant branch star of mass equal to 0.88~\\msun \\ and a companion star of mass ranging from 0.9 down to 0.1~\\msun. We first compare the results obtained using two different numerical techniques with different resolutions, and find overall very good agreement. We then compare the outcomes of those simulations with observed systems thought to have gone through a common envelope. The simulations fail to reproduce those systems in the sense that most of the envelope of the donor remains bound at the end of the simulations and the final orbital separations between the donor's remnant and the companion, ranging from 26.8 down to 5.9~\\rsun, are larger than the ones observed. We suggest that this discrepancy vouches for recombination playing an essential role in the ejection of the envelope and/or significant shrinkage of the orbit happening in the subsequent phase. ",
        "introduction": "\\label{sec:intro}  Around 60\\% of F and G stars are binaries, of which about 30\\% have separations smaller than 30 AU and will interact during the primary's evolution \\citep{Duquennoy1991}. During the giant phases of the primary, companions closer than $\\sim 5$ AU enter a strong interaction phase with the primary and, under certain circumstances, a common envelope (CE) may form around the two stars. The secondary star spirals inside the envelope of the primary and may also fill its own Roche lobe because it cannot accrete all the matter coming from the donor star. This process is called a \\textit{common envelope interaction} and was originally described by \\cite{Paczynski1976}. For a general review of the topic see, e.g., \\cite{Iben1993}. There are two different processes leading to the onset of a CE phase: the start of unstable mass transfer from the expanding primary to the secondary \\citep{HjellmingWebbink1987, HurleyEtAl2002} and the development of a tidal instability that occurs if there is not enough angular momentum in the orbit to maintain the primary's envelope in synchronization \\citep{Darwin1879}. The post-CE system will be either a compact binary system, if there is enough energy to eject the primary's envelope, or a merger, if not.  The CE interaction is an essential ingredient for any binary population synthesis study of intermediate \\citep[e.g.,][]{PolitanoEtAl2010} or massive stars \\citep[e.g.,][]{Belczynski2008}. Compact binaries are believed to be formed through at least one CE phase. Among them are symbiotic binaries, supersoft X-ray sources, cataclysmic variables and double white dwarfs, which are all possible supernova Type Ia progenitors. As \\cite{MengEtAl2010} pointed out, results deduced from population synthesis studies such as the Type Ia supernova birth rate are highly dependent on the physics of the CE phase. Therefore, it is paramount to understand more accurately the CE interaction in order to identify the formation channels of such supernovae and to compare observations with predictive models. Moreover, many substellar companions to evolved stars have recently been discovered with small orbital separation. \\cite{Maxted2006} found a brown dwarf orbiting a white dwarf with a 116 min period, while \\cite{SetiawanEtAl2010} discovered a system composed of a Jupiter-like object orbiting an horizontal branch star with a 16.2 days period. We therefore know that substellar companions can survive a CE interaction, but what is the minimum mass of the companion that can eject the envelope? Is the ejected envelope entirely unbound or will some of it eventually fall back and form a circumbinary disk? Were the substellar companions present before and survived the CE or were they formed later on in such a disk \\citep{Perets2010}? Those questions remain unanswered.  Although the CE process was outlined more than 30 years ago, it is still far from understood quantitatively. Numerical simulations suggest that the typical duration of the entire CE phase is short --- less than $10^3$ years --- which makes CE ejections unlikely to be observed. However, one can use observations of post-CE binaries to better understand CE evolution. With the use of stellar models, the initial configuration of such systems can be approximately determined from the final configuration. Using either the \\ace-formalism \\citep[][but see \\citealt{AlphaPaper2011}]{Webbink1984} or the $\\gamma$-formalism \\citep{Nelemans2000} the relevant parameters can be constrained and the CE ejection efficiency can be predicted. Using this approach, \\cite{AlphaPaper2011} suggested an anti-correlation between $\\alpha$, the CE efficiency parameter, and the secondary to primary mass ratio.  The entire CE evolution can be divided into three different phases \\citep{Podsiadlowski2001} with different timescales, length scales and physics involved. These differences are the reasons why reproducing the entire CE evolution of a given system accurately is challenging. Therefore, one usually treats one phase after the other with different methods. In this paper, we focus only on the rapid infall phase, which has a short timescale ($\\sim 1-10$ years), and in which the evolution is driven by drag forces. Several numerical hydrodynamic studies of the CE interaction have been carried out in the past \\citep[for an exhaustive list, see][]{TaamSandquist2000}, including a series of ten papers starting with the two-dimensional calculation of the interaction of a 16~\\msun \\ supergiant and a 1~\\msun \\ neutron star \\citep{PaperII}, and most recently treating three-dimensional simulations of the CE interaction between 3 or 5~\\msun \\ giant stars and 0.4 or 0.6~\\msun \\ main sequence (MS) companions \\citep{Sandquist1998}. The latter study has been extended first by \\cite{SandquistEtAl2000} to 1~\\msun \\ and 2~\\msun \\ red giant branch (RGB) stars with companion masses ranging from 0.1 to 0.45~\\msun, then by \\cite{DeMarco2003} to a 1~\\msun \\ asymptotic giant branch star with a 0.1 or 0.2~\\msun \\ companion. \\cite{RickerTaam2008} computed high resolution simulations of the CE phase between a 1.05~\\msun \\ RGB star and a 0.6~\\msun \\ compact companion, and concluded that the gravitational component of the drag dominates over the hydrodynamical component \\citep[also see][]{TaamRicker2010, RickerTaam2011}.  A direct comparison of the results obtained using different numerical methods has however never been carried out. Although analytical/empirical work has included discussion regarding observational data, there are only a couple of publications that connect simulations and observations in a meaningful way \\citep[see e.g.,][]{SandquistEtAl2000}. Those are, as we will explain in \\S\\ref{sec:code} and \\S\\ref{sec:results}, key steps to better understand the implications of CE interactions and the physical processes driving them. In this paper we therefore present numerical simulations with two different algorithms of the CE interaction of a 0.88~\\msun \\ RGB star with a MS companion. Different companion masses from 0.1~\\msun \\ to 0.9~\\msun \\ are considered. The simulations are carried out with both an Eulerian code ({\\it Enzo} in uniform-grid mode,  \\cite{OsheaEtAl2004} and {\\it enzo.googlecode.com}) and a Lagrangian code ({\\it SNSPH}, \\citealt{FryerEtAl2006}), and for different resolutions. We describe the numerical methods and the initial conditions of our 15 simulations in \\S\\ref{sec:code} and \\S\\ref{sec:simus}. We describe and discuss the results in \\S\\ref{sec:results} and \\S\\ref{sec:discussion}, and finally conclude and summarize in \\S\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Comparison of simulations and observations} \\label{subsec:alphacomp}  We now compare the numerical results with a sample of 61 observed post-CE systems listed in \\cite{Zorotovic2010} and \\cite{AlphaPaper2011}.  \\subsubsection{Final separations}  For a given companion mass (or alternatively mass ratio $q$) we obtain 3 values for the final separation $A_f$, one for each simulation carried out with that companion mass (Table \\ref{tab:runs} and Fig.~\\ref{fig:sepsims}). One can distinguish between these values at high $q$ ($q \\geq 0.34$), which correspond to ``heavy'' companions ($M_2 \\geq 0.3$~\\msun), and the ones at low $q$ ($q < 0.34$) corresponding to ``light'' companions ($M_2 < 0.3$~\\msun). At high $q$, the values of $A_f$ are very similar and the standard deviation is more than 20 times smaller than the average value of $A_f$. At low $q$, the companion sinks deeper and as a consequence, the resolution used in the $128^3$ {\\it Enzo} simulations is not sufficient. However, as one increases the resolution to $256^3$ cells, the final separations converge to the solutions given by the {\\it SNSPH} simulations.  \\begin{figure}[h!] \\begin{center} \\includegraphics[scale=0.45]{sepsims.eps} \\caption{Final separations as a function of the mass ratio $q$ for the {\\it SNSPH} (black cross), {\\it Enzo} $128^3$ (blue circle) and {\\it Enzo} $256^3$ (red triangle) simulations. \\label{fig:sepsims} } \\end{center} \\end{figure}  Fig.~\\ref{fig:histo} shows the distribution of orbital separations reached by the 61 post-CE systems. For all these systems, there has been no substantial orbital shrinkage due to phenomena such as magnetic braking or radiation of gravitational waves \\citep[see discussion in][]{Schreiber2003}. Although they cover a significant range in secondary masses, going from a 1.1~\\msun \\ MS star down to a 0.05~\\msun \\ brown dwarf, all of them have separations smaller than 11~\\rsun. Furthermore, 87\\% of those systems have separations smaller than 4~\\rsun, which is smaller than any value obtained in our simulations. This is even more obviously shown in Fig.~\\ref{fig:sepobs}, where the final separations for simulations presented here and in the literature are compared to the orbital separations of the observed post-CE systems. Although a couple of observed systems have $q \\geq 0.5$, one clearly sees that the simulations with $M_2 = 0.9$ and 0.6~\\msun \\ leave the companion far out. Systems with lower mass companions ($M_2 \\leq 0.3$~\\msun) have by and large lower orbital separations than in our simulations. Simulations of \\cite{Sandquist1998} and \\cite{RickerTaam2008} shown in Fig.~\\ref{fig:sepobs} give results consistent with ours. All these numerical simulations suggest that the separations between the secondary and the primary's remnant at the end of the simulated rapid infall phase are too large to explain the orbital separation of the currently observed post-CE systems. This suggests that further evolution of the orbital separation must occur during the phase immediately following the rapid infall phase. We discuss this point further in \\S\\ref{subsec:reasons}.  \\begin{figure}[h!] \\begin{center} \\includegraphics[scale=0.5]{SystemsHisto.eps} \\caption{Distribution of post-CE systems as a function of their observed orbital separation from \\cite{Zorotovic2010} and \\cite{AlphaPaper2011}. \\label{fig:histo} } \\end{center} \\end{figure}  \\begin{figure}[h!] \\begin{center} \\includegraphics[scale=0.39]{sepobs.eps} \\caption{Comparison between the orbital separations of observed post-CE systems (black dot) and the final separations reached at the end of the simulations (red circle), as well as the ones by \\cite{Sandquist1998} (green circles) and by \\cite{RickerTaam2008} (blue triangle). \\label{fig:sepobs} } \\end{center} \\end{figure}  \\subsubsection{The state of the envelope at the end of the simulations} \\label{sec:envelope}  As shown on Table~\\ref{tab:masses}, most of the primary's envelope remains bound in all of our simulations. We study the situation in detail for our canonical model with the 0.6~\\msun\\ companion here. The evolution of the mass for different components is plotted in Fig.~\\ref{fig:massbound}. It first confirms that some envelope mass is unbound only during the first 50~days, after which neither angular momentum (Fig.~\\ref{fig:AM}) nor kinetic energy (Fig.~\\ref{fig:energy}) are exchanged between the unbound mass and the rest of the system. It also shows that more than 85 \\% of the mass remains bound at the end of the simulation. This outcome, already pointed out by \\cite{Sandquist1998}, is quite intriguing, since the post-CE binaries observed must have succeeded in ejecting their envelope. After about 400 days, most of the envelope mass in our models has been moved to a larger radius ($\\sim 100$~\\rsun, see bottom panel in Fig.~\\ref{fig:vescape}), well outside the orbit of the primary core and the companion but remaining bound.  \\begin{table}[h!] \\begin{center} \\scalebox{0.9} {\\begin{tabular}{ccc} \\hline \\hline & $M_2$ (\\msun) & $M_{\\rm bound}$ (\\msun)\\tablenotemark{a}\\\\ \\hline SPH1 & 0.9 & 0.44 \\\\ SPH2 & 0.6 & 0.44 \\\\ SPH3 & 0.3 & 0.45 \\\\ SPH4 & 0.15 & 0.46 \\\\ SPH5 & 0.1 & 0.48 \\\\ \\hline \\hline \\end{tabular}} \\caption{Amount of the envelope mass still bound at the end of the {\\it SNSPH} simulations. \\tablenotetext{a}{At the start of the simulations, $M_{\\rm bound}$ equals the total envelope mass $M_e \\equiv M_1 - M_c = 0.49$~\\msun} \\label{tab:masses} } \\end{center} \\end{table}  \\begin{figure}[h!] \\begin{center} \\includegraphics[scale=0.45]{M_SPH_WD06.eps} \\caption{Evolution of the total mass (solid line), the bound mass (dashed line), the mass within the volume of the {\\it Enzo} grid (dotted line), the mass within the initial volume of the primary (dash-cross line) and the mass within the orbital separation (dash-dot line) for the SPH2 simulation. \\label{fig:massbound} } \\end{center} \\end{figure}   \\begin{figure}[h!] \\begin{center} \\includegraphics[scale=0.5]{vescape.eps} \\caption{Top: Comparison between the escape velocity (dashed line) and the radial velocity (black dots) of the final system for the SPH2 simulation. Bottom: Mass enclosed as a function of radius. \\label{fig:vescape} } \\end{center} \\end{figure}  We now investigate how bound the final system is. We consider the center of mass of the system composed by the secondary and the mass within the current orbit as the center of our frame of reference. Then, we partition the domain into concentric shells with identical thickness, calculate the average radial velocity of each shell and compare it to the escape velocity at that location. Fig.~\\ref{fig:vescape} shows the escape velocity and the average radial velocity of the shells. The radial velocity is always positive and is similar to the space velocity at radii larger than 600 \\rsun, as expected for envelope ejection. At radii smaller than 300 \\rsun, the radial velocity is much smaller than the space velocity, suggesting that orbital motions dominate at those radii. All the mass within $10^3$ \\rsun \\ is bound, which corresponds to more than 85\\% of the envelope mass. The remaining mass is found at radii between $10^3$ and $6 \\times 10^3$\\rsun, where the radial velocities are typically between 25 and 75 \\kms. Those particles were initially in the outer parts of the giant star, and were the first to encounter the secondary. At that time of the in-spiral, the shock was slightly supersonic ($V_2 \\sim 35$~\\kms \\ and $c_s \\sim 20$~\\kms). This regime of evolution is thus different from later phases when the secondary sinks deeper into the primary's envelope, where its velocity does not really increase (Fig.~\\ref{fig:velocity}) but the sound speed of the medium does (Fig.~\\ref{fig:profiles}).  We can measure how much extra energy would be required to unbind the envelope at each radius, using the definition  \\begin{equation} E_{\\rm extra} = \\sum_i \\frac{1}{2}  M_i  \\left(v_{e,i} - v_{r,i}  \\right)^2 \\end{equation}  \\noindent where $v_{e,i}$ and $v_{r,i}$ are the escape velocity at the location of the $i$-th shell and its average radial velocity, respectively. One finds $E_{\\rm extra} \\sim 8.4 \\times 10^{45}$ ergs which represents just over 10\\% of the initial binding energy of the primary envelope. Thus, a relatively small additional input of energy could be sufficient to completely unbind the remaining envelope material.   We have compared here the final separations deduced from observations and those determined from the simulations. We have purposefully stayed away from calculating the ejection efficiency \\ace \\ \\citep{Webbink1984, AlphaPaper2011}. Indeed, we question what the relevance of calculating \\ace \\ is when the envelope has not yet been fully ejected, true both in the \\cite{Sandquist1998} and our simulations.  We therefore defer for the moment the task of calculating \\ace \\ from simulation --- a long term goal of this project --- until the simulations are more advanced.  In conclusion, the hydrodynamic simulations do not reproduce the post-CE systems in the sense that the system is left at too large separations and the envelope is not unbound at the end of the rapid infall phase. This means that either physical processes that are not accounted for in the simulations are responsible for the envelope ejection, or the envelope ejection and a significant reduction of the orbit actually happens during the later subsequent slow in-spiral phase. We discuss both possibilities in the following section.  \\subsection{Reproducing the observations} \\label{subsec:reasons}  In this section we first study and quantify physical processes that are not taken into account in our hydrodynamic simulations and that might be responsible for ejecting the envelope. Then, we focus on the subsequent slow in-spiral phase and investigate whether the envelope can be ejected and the separation significantly reduced during this subsequent phase.  \\subsubsection{Rotation of the primary} \\label{subsub:rotation}  The envelope of the progenitor is initially non-rotating and although the calculation done in \\S\\ref{sec:simus} shows that, regardless of the initial rotation velocity of the envelope, its rotational energy is negligible in comparison with its binding energy, we suspected at first that the absence of rotation might be the reason for most of the envelope to remain bound. However, \\cite{Sandquist1998} carried out two identical simulations where they modified the initial rotation state of the primary from a giant star in synchronization with the orbit to a non-rotating one (their simulations 1 and 2). In both cases, the evolution of the bound mass and the final orbital parameters are similar.  It thus does not seem that changing the initial rotation of the primary leads to a different CE outcome.  \\subsubsection{Physics not included in the simulations} \\label{subsub:missing}  The hydrodynamics codes use an ideal gas equation of state (\\S\\ref{subsec:codesdiff}) which, by definition, does not include variable abundances and the different ionization layers of the envelope. \\cite{Han1995} suggested that recombination might play a role in CE interactions. As the outer parts of the envelope expand and cool, ions recombine with electrons, releasing energy that could aid in unbinding the envelope. Although it is unclear how efficient this process is and how much of the initial recombination budget can be used, one can calculate an upper limit on how much energy can be injected into the envelope by recombination.  According to our stellar evolution model, the hydrogen fraction within the convective envelope of our RGB star is $X \\sim 0.68$. The mass of the envelope is $M_e = 0.49$~\\msun \\ and each proton recombining with an electron produces an energy $E_0 = 13.6$~eV. We also have to calculate how much of the envelope is ionized. Therefore, we calculate the partition functions $Z$ for hydrogen. The hydrogen ion has no degeneracy so $Z_2 = 1$. The partition function for the hydrogen atom at temperature $T$ is  \\begin{equation} Z_1 = \\sum_{n=1}^{\\infty} 2n^2 \\exp{\\frac{E_0(1/n^2-1)}{kT}} \\label{eq:partition} \\end{equation}  \\noindent where $k=8.6173 \\times 10^{-5}$~eV.$K^{-1}$ is the Boltzmann constant. We truncate the sum in Eq.~\\ref{eq:partition} at the first integer $n_{\\rm max}$ such that the distance at which the electron orbits the proton for this quantum number is larger than $l_{\\rm max}=10^{-6}$~cm, i.e. $a_0n_{\\rm max}^2 > l_{\\rm max}$ where $a_0 = 5.2918 \\times 10^9$~cm is the Bohr radius \\citep{Miranda2001}. We then use the Saha formula to calculate the ratio of ionized to neutral hydrogen \\citep{Carroll}:  \\begin{equation} N_2/N_1 = \\frac{2Z_2}{n_e Z_1}\\left(\\frac{2\\pi m_e k T}{h^2}\\right)^{3/2} \\exp{(-E_0/kT)} \\end{equation}  \\noindent where $n_e$ is the number density of free electrons and $m_e$ is the electron mass. We find that 91\\% of the envelope is ionized. Consequently, the recombination of the whole ionized envelope would produce an extra energy  \\begin{equation} E_{\\rm recomb} = 0.91 \\times XM_e \\frac{N_A}{M_H} \\times 13.6~{\\rm eV} \\label{eq:recomb} \\end{equation}  \\noindent where $N_A$ is the Avogadro number and $M_H$ the atomic mass of hydrogen. One finds $E_{\\rm recomb} \\sim 1.18 \\times 10^{46}$~ergs, which is slightly higher than the extra energy $E_{\\rm extra}$ required to eject the envelope in our canonical model (\\S\\ref{sec:envelope}). Thus, we conclude that recombination in the envelope could substantially aid in unbinding it.   Another source of energy could be radiation pressure. For low- and intermediate-mass giants in hydrostatic equilibrium, radiation pressure ($P_{rad} \\equiv a T^4/3$, where $a$ is the radiation constant) is negligible compared to gas pressure (Eq.~\\ref{eq:hd4}): for our primary, $P_{rad} / P_{gas} \\la 0.01$ except in a small zone ($0.1$ \\rsun $ \\leq r \\leq 10$ \\rsun), where $P_{rad} / P_{gas} \\la 0.1$. However, the deep in-spiral of the companion within the primary's envelope will induce local shock heating. The increase of temperature is proportional to the square of the Mach number \\citep{TarbellEtAl1999}, so even if the companion is orbiting at twice the local sound speed, the radiation pressure to  gas pressure after the shock becomes:  \\begin{equation} \\left( \\frac{P_{rad}}{P_{gas}} \\right)^{\\rm after} \\propto \\left( \\frac{P_{rad}}{P_{gas}} \\right)^{\\rm before} (M^2)^3 = 6.4 \\end{equation}  \\noindent Therefore, including radiation pressure in the equation of state will increase the total pressure locally and might reduce the energy required to eject the envelope. However, it is possible that this effect is globally small, since this extra heating source is probably very localized around the companion.  \\subsubsection{The post-rapid-infall phase}  At the end of the rapid infall phase, the orbit is stable until the end of the simulations (a few more years). Consequently, there is no further hydrodynamical coupling between the extended envelope and the surviving binary. We now investigate whether the envelope is likely to be ejected during this slower in-spiral phase.  Although the resolution of the simulations prevents us from quantifying how much envelope will be left around the core of the primary, one can still describe qualitatively what the evolution of the primary's remnant will be. Fig.~\\ref{fig:massbound} shows that less than $10^{-2}$~\\msun \\ is left around the primary's core, so the primary will depart the giant branch (\\citealt{Bloecker1995}, but see also the discussion in \\citealt{AlphaPaper2011}). Then two scenarios might occur depending on how long the partially ejected envelope will take to fall back.  If the star is given enough time to transit to the blue due to hydrogen burning at the base of the envelope before the lifted envelope falls back, the star will readjust on its thermal timescale of the remaining envelope, and eventually end its life as a Helium white dwarf. This transition will last $\\sim 10^3$~years during which the star will have a luminosity between 300 and 1\\,000 \\lsun \\ \\citep[][their Fig.~1]{IbenTutukov1993}, which is consistent with the more recent work of \\cite{DriebeEtAl1998} (their Fig.~1). If we assume the remnant to have a luminosity $L_c \\sim 500$~\\lsun, we can compare the gravitational acceleration of a gas particle with the radiation acceleration defined by  \\begin{equation} a_{rad} = \\frac{L_c}{4\\pi r^2} \\frac{\\kappa}{c} \\label{eq:radacc} \\end{equation}  \\noindent where $r$ is the distance between the gas particle and the core, $\\kappa = 0.4$~cm$^2$\\,$g^{-1}$ is the opacity for Thompson scattering for hydrogen, and $c$ is the speed of light. We still find the radiation acceleration to be overall almost two orders of magnitude smaller than the gravitational acceleration.  If on the contrary, the envelope falls back before the primary's remnant had crossed the Hertzsprung-Russell diagram, a circumbinary disk will form \\citep{KashiSoker2011}. They refer to the numerical work done by \\cite{ArtymowiczEtAl1991}, which suggests that in such a configuration, the binary separation will decrease due to Lindblad resonances --- mainly --- as well as viscous tides. Although this mechanism has the advantage of explaining how the orbital separation will diminish during the subsequent phase, the ability of radiation to eject the gas will even be reduced in comparison with the previous situation, so it is not clear how the latter will eventually be unbound.  In conclusion, radiation acceleration alone does not seem to be responsible for unbinding the remaining gas, regardless of the time the partially ejected envelope will remain suspended for."
    },
    "1107/1107.4074_arXiv.txt": {
        "abstract": "We have analyzed $^{137}$Cs decay data, obtained from a small sample onboard the MESSENGER spacecraft en route to Mercury, with the aim of setting limits on a possible correlation between nuclear decay rates and solar activity.  Such a correlation has been suggested recently on the basis of data from $^{54}$Mn decay during the solar flare of 13 December 2006, and by indications of an annual and other periodic variations in the decay rates of $^{32}$Si, $^{36}$Cl, and $^{226}$Ra. Data from five measurements of the $^{137}$Cs count rate over a period of approximately 5.4 years have been fit to a formula which accounts for the usual exponential decrease in count rate over time, along with the addition of a theoretical solar contribution varying with MESSENGER-Sun separation. The indication of solar influence is then characterized by a non-zero value of the calculated parameter $\\xi$, and we find $\\xi=(2.8\\pm8.1)\\times10^{-3}$ for $^{137}$Cs. A simulation of the increased data that can hypothetically be expected following Mercury orbit insertion on 18 March 2011 suggests that the anticipated improvement in the determination of $\\xi$ could reveal a non-zero value of $\\xi$ if present at a level consistent with other data. ",
        "introduction": "The presence of the $^{137}$Cs source onboard MESSENGER provides an opportunity to test whether the $^{137}$Cs decay rate varies with MESSENGER-Sun distance. Following the discussion presented in \\citet{fis09} we assume that the beta decay rate $dN(t)/dt \\equiv \\dot{N}(t)$ of a sample containing $N(t)$ unstable nuclei can be written in the form \\begin{eqnarray} \\frac{-dN(t)}{dt} = \\lambda(t)N(t) = [\\lambda_0 + \\lambda_1(\\vec{r},t)]N(t), \\label{Decay} \\end{eqnarray} \\noindent{}where $\\lambda_0$ represents the intrinsic contribution to the $\\beta$-decay rate arising from the conventional weak interaction along with a possible time-independent background arising from new interactions. The time-dependent perturbation $\\lambda_1(\\vec{r},t)$ is also a function of the distance $r = |\\vec{r}|$ between the decaying nucleus and the source, here assumed to be the Sun.  For a perturbation whose influence on a decaying nucleus varies as $1/r^n$, where $n$ is an integer, $\\lambda_1(\\vec{r},t)$ can be expressed in the form \\begin{eqnarray} \\lambda_1(\\vec{r},t) = \\lambda_0\\xi^{(n)} \\left[ \\frac{R}{r(t)} \\right]^n , \\label{lambda1} \\end{eqnarray} \\noindent where $R \\equiv r(t=0)$, and $\\xi^{(n)}$ is the parameter we are interested in constraining. In what follows we focus initially on $\\xi^{(2)}\\equiv\\xi$, which is the most likely variation that we expect, (i.e., an inverse-square law). Results for $n=1$ and $n=3$, which are the next most likely alternatives to $n=2$, can be analyzed in a similar way. We note from Eqs. (\\ref{Decay}) and (\\ref{lambda1}) that $\\xi$ is not a universal constant since it depends on the specified initial values $R$ and $t=0$.  In the present paper we define $t=0$ as the launch time, in which case $R = 1~\\textrm{AU} = 1.495979 \\times 10^{8}$ km. Integrating Eq. (\\ref{Decay}) we find \\begin{eqnarray} N(t) = N_0 \\exp \\{ -\\lambda_0 [t + \\xi I(t)]\\}, \\label{Noft} \\end{eqnarray} \\begin{eqnarray} \\nonumber  -\\frac{dN(t)}{dt} = \\\\ \\lambda_0N_0 \\exp\\left(-\\lambda_0t\\right) \\left\\lbrace 1 + \\xi \\left[ \\frac{R^2}{r^2(t)} - \\lambda_0 I(t)\\right]\\right\\rbrace, \\label{fulldecay} \\end{eqnarray} \\begin{eqnarray} I(t) = \\int_0^t dt^\\prime \\left[ \\frac{R}{r(t')}\\right]^2. \\label{I2} \\end{eqnarray} \\noindent We note from Eqs. (\\ref{Noft}-\\ref{I2}) that since $I(0) = 0$, the effective $^{137}$Cs decay constant at $t = 0$ is $\\lambda \\equiv \\lambda_0(1 + \\xi)$, where $\\lambda = 6.311(19) \\times 10^{-5}~\\textrm{d}^{-1}$ is the standard value measured on Earth (at 1 AU) corresponding to $T_{1/2} = 30.07(9)$ yr. It is convenient to rewrite Eq. (\\ref{fulldecay}) in terms of directly measured quantities by first expressing $\\lambda_0$ in terms of $\\lambda$, and then eliminating the unknown $N_0$ by taking appropriate ratios. To lowest order in $\\xi \\ll 1$ we find \\begin{eqnarray} \\nonumber  e^{+\\lambda t} \\frac{dN(t)/dt}{dN(t=0)/dt} &\\cong& 1 + \\xi \\left[\\frac{R^2}{r^2(t)} - \\lambda I(t) + \\lambda t-1\\right] \\\\ &\\equiv& 1 + \\xi B(t). \\label{Boft} \\end{eqnarray} \\noindent We note from Eq. (\\ref{Boft}) that since $B(t=0) = 0$ the count rate ratio approaches unity as $t \\rightarrow 0$, as expected. Comparing Eqs. (\\ref{fulldecay}) and (\\ref{Boft}), the additional terms in the square brackets in Eq. (\\ref{Boft}) arise from computing the indicated ratio, and then expressing the final result in terms of the laboratory value $\\lambda$.  In principle, the left hand side of Eq. (\\ref{Boft}) and $B(t)$ are directly measurable, and hence can be used to set limits on $\\xi$. These limits can then be compared to those obtained in the laboratory.  In practice, however, $dN(t=0)/dt$ was not measured, and hence our limits on $\\xi$ were determined from ratios of count rates obtained from measurements made during the five cruise periods. A summary of the data obtained from these measurements is presented in Table \\ref{Tbl:Detail}. Among the entries in this table, the precision of the data which enter into the calculation of $B(t)$ is such that the errors on $B(t)$ can be neglected compared to those arising from the count rates (see discussion below). The integral $I(t)$, which is exhibited in Figure \\ref{fig:Rcurves}, accounts for the cumulative $\\xi$-dependent contributions to the $^{137}$Cs decay rate arising from both the change in MESSENGER-Sun distance, as well as from the normalization relation $\\lambda = \\lambda_0(1 + \\xi)$.  \\begin{table*} \\footnotesize \\caption{Summary of $^{137}$Cs decay data from the MESSENGER mission. For each of the five measurement periods ($t_0,t_1,t_2,t_3,t_4$), these data include the calculated values of $I(t)$, $B(t)$, and $\\Delta(t)$, for a perturbation varying as $1/r^2$\\label{Tbl:Detail}} \\begin{tabular}{@{}lrrrrr@{}} \\tableline & Cruise 0 & Cruise 1 & Cruise 2 & Cruise 3 & Cruise 4 \\\\ &  Initial checkout & Mercury Flyby 1 & Mercury Flyby 2 & Mercury Flyby 3 & Final checkout \\\\ \\tableline Calendar date & 16 Nov 2004 & 14 Jan 2008 & 6 Oct 2008 & 2 Oct 2009 & 14 Apr 2010 \\\\ Julian date  &  2453322.500  &  2454478.500  &  2454744.500  &  2455102.500 & 2455300.500 \\\\ Distance from the Sun (km)  &  160599579.0277  &  54595403.8012  &  52751315.2631  &  48132083.3548 & 63523375.7902 \\\\ Elapsed time since Cruise 0 (days)  &  0.0  &  1156.0  &  1422.0  &  1780.0 & 1978.0 \\\\ Mean count rate in peak (cps) &  0.100$\\pm$0.0012  &  0.0942$\\pm$0.0019  &  0.0933$\\pm$0.0020  &  0.0886$\\pm$0.0017 & 0.0886$\\pm$0.0020\\\\ Total accumulation time (hours) & 60.00 &  45.00  &  40.00  &  60.00  &  39.00 \\\\ $I(t)$ (days)  &  91.173  &  2284.190  &  3408.077  &  5162.823 & 6228.370  \\\\ $B(t)$  &  -0.1316  &  6.4397  &  6.9235  &  8.4530 & 4.2842 \\\\ $\\Delta(t)$  &  0.0000  &  6.5751  &  7.0551  &  8.5847 & 4.4159 \\\\ \\tableline \\end{tabular}  \\end{table*}      \\begin{figure}[h] \\includegraphics[width=\\columnwidth]{fig2.eps} \\caption{Plots of the integrals $I(t)$, for $n=1,2,3,$ in Eq. (\\ref{lambda1}).  The integrals are defined in analogy to $I^{(2)}(t)$ in Eq. (\\ref{I2}). \\label{fig:Rcurves}} \\end{figure}   Returning to Eq. (\\ref{Boft}), the experimentally interesting quantities are the ratios \\begin{eqnarray} \\nonumber \\frac{[\\exp(\\lambda t_i)dN(t_i)/dt]}{[\\exp (\\lambda t_0)dN(t_0)/dt]}, \\end{eqnarray} \\noindent where $t_0$ and $t_i$ $(i = 1,2,3,4)$ denote the start times of the five cruise measurement periods.  These five measurements thus determine four ratios which, along with the corresponding values of $B(t_i$), can be used to obtain $\\xi$. ",
        "conclusions": ""
    },
    "1107/1107.3148_arXiv.txt": {
        "abstract": "Heavily obscured ($N_{\\rm H}\\ga3\\times10^{23}$~cm$^{-2}$) Active Galactic Nuclei (AGNs) not detected in even the deepest X-ray surveys are often considered to be comparably numerous to the unobscured and moderately obscured AGNs. Such sources are required to fit the cosmic X-ray background (XRB) emission in the 10--30 keV band. We identify a numerically significant population of heavily obscured AGNs at $z\\approx0.5$--1 in the \\chandra\\ Deep \\hbox{Field-South} (CDF-S) and Extended \\chandra\\ Deep Field-South by selecting 242 X-ray undetected objects with infrared-based star formation rates (SFRs) substantially higher (a factor of 3.2 or more) than their SFRs determined from the UV after correcting for dust extinction. An \\hbox{X-ray} stacking analysis of 23 candidates in the central CDF-S region using the 4 Ms \\chandra\\ data reveals a hard X-ray signal with an effective power-law photon index of $\\Gamma=0.6_{-0.4}^{+0.3}$, indicating a significant contribution from obscured AGNs. Based on Monte Carlo simulations, we conclude that $74\\pm25\\%$ of the selected galaxies host obscured AGNs, within which $\\approx95\\%$ are heavily obscured and $\\approx80\\%$ are Compton-thick (CT; $N_{\\rm H}>1.5\\times10^{24}$~cm$^{-2}$). The heavily obscured objects in our sample are of moderate intrinsic X-ray luminosity [$\\approx(0.9$--$4)\\times10^{42}$~\\lum\\ in the 2--10 keV band]. The space density of the CT AGNs is $(1.6\\pm0.5)\\times10^{-4}$~Mpc$^{-3}$. The $z\\approx0.5$--1 CT objects studied here are expected to contribute $\\approx1\\%$ of the total XRB flux in the 10--30 keV band, and they account for $\\approx5$--15\\% of the emission in this energy band expected from all CT AGNs according to population-synthesis models. In the 6--8 keV band, the stacked signal of the 23 heavily obscured candidates accounts for $<5\\%$ of the unresolved XRB flux, while the unresolved $\\approx25\\%$ of the XRB in this band can probably be explained by a stacking analysis of the X-ray undetected optical galaxies in the CDF-S (a 2.5~$\\sigma$ stacked signal). We discuss prospects to identify such heavily obscured objects using future hard X-ray observatories. ",
        "introduction": "Deep X-ray surveys have provided the most effective method of identifying reliable and fairly complete samples of Active Galactic Nuclei (AGNs) out to $z\\approx5$ (e.g., see \\citealt{Brandt2005} for a review). The observed AGN sky density reaches $\\approx10\\,000$~deg$^{-2}$ in the deepest X-ray surveys, the \\chandra\\ Deep Fields \\citep[e.g.,][]{Bauer2004,Xue2011}. These \\hbox{X-ray} point sources are largely responsible for the observed cosmic X-ray background (XRB). A significant portion ($\\approx70\\textrm{--}90\\%$) of the XRB in the \\hbox{0.5--6} keV range has been resolved into discrete sources by \\chandra\\ and {\\it XMM-Newton} \\citep[e.g.,][]{Moretti2002,Bauer2004,Hickox2006}, the majority of which are AGNs with moderate X-ray obscuration ($N_{\\rm H}\\la3\\times10^{23}$~cm$^{-2}$; e.g., \\citealt{Szokoly2004,Barger2005,Tozzi2006}). However, the resolved fraction of the XRB decreases toward higher energies, being $\\approx60\\%$ in the 6--8 keV band and $\\approx50\\%$ in the 8--12~keV band \\citep[][]{Worsley2004,Worsley2005}. Population-synthesis models suggest that at $\\approx20$--30~keV, where the XRB reaches its peak flux \\citep[e.g.,][]{Gruber1999,Moretti2009}, the unobscured or moderately obscured AGNs discovered at lower energies cannot account for the XRB flux entirely, and an additional population of heavily obscured ($N_{\\rm H}\\ga3\\times10^{23}$~cm$^{-2}$), or even Compton-thick ($N_{\\rm H}>1.5\\times10^{24}$~cm$^{-2}$, hereafter CT), AGNs at mainly $z\\approx0$--1.5 is required. The number density of these heavily obscured AGNs is estimated to be of the same order as that of moderately obscured AGNs \\citep[e.g.,][]{Gilli2007}. However, as the X-ray emission (below at least 10 keV) of such sources is significantly suppressed, only a few distant heavily obscured AGNs have been clearly identified even in the deepest X-ray surveys \\citep[e.g.,][]{Tozzi2006,Alexander2008,Georgantopoulos2009,Comastri2011,Feruglio2011}, and thus a significant fraction of the AGN population probably remains undetected.  Infrared (IR) selection is a powerful tool for detecting heavily obscured AGNs that cannot be identified in the X-ray band. X-ray and UV emission absorbed by the obscuring material is reprocessed and reemitted mainly at mid-IR wavelengths \\citep[e.g.,][]{Silva2004,Prieto2010}. The mid-IR emission is less affected by dust extinction than at optical or near-IR wavelengths, rendering it more suitable for identifying obscured AGNs. With the {\\it Spitzer} mission \\citep{Werner2004}, deep mid-IR data in multiple bands can be obtained by the Infrared Array Camera (IRAC) and Multiband Imaging Photometer (MIPS) instruments.  Various IR selection methods for obscured AGNs have been proposed. {\\it Spitzer} IRAC power-law selection chooses sources whose IRAC spectral energy distributions (SEDs) follow a power law with a slope of $\\alpha\\le-0.5$ ($f_{\\nu}\\propto\\nu^{\\alpha}$; e.g., \\citealt{Alonso2006,Donley2007,Donley2010}). This criterion was determined based on the average optical-to-IR spectral properties of X-ray selected QSOs \\citep[e.g.,][]{Elvis1994}. It can select a relatively pure sample of AGNs with little contamination from star-forming galaxies but is generally limited to the most-luminous sources \\citep[e.g.,][]{Polletta2008}. IRAC color-color selection applies color cuts in the IRAC color space to identify AGNs \\citep[e.g.,][]{Lacy2004,Stern2005}. This method is based on the IRAC color properties of optically selected QSOs and AGNs, and it tends to have significant contamination from star-forming galaxies when applied to deep mid-IR data \\citep[e.g.,][]{Cardamone2008,Georgantopoulos2008,Brusa2010}. IR-excess selection uses a combined UV--optical--IR color cut to identify AGNs \\citep[e.g.,][]{Dey2008,Fiore2008,Lanzuisi2009}. This technique is also contaminated by star-forming galaxies \\citep[e.g.,][]{Donley2008}.   Compared to the above methods, which utilize observational data directly and determine the selection criteria empirically, \\citet{Daddi2007} adopted a generally different approach to find obscured AGNs. They selected candidates with significantly higher IR-based (also including the component corresponding to the transmitted UV emission; see \\S2.1 below for details) star formation rates (SFRs) than their UV-based (after correcting for dust extinction) SFRs. The amount of excess in the IR-based SFR is a measurement of the excess in the IR emission, which is likely caused by reprocessed emission from the dusty torus of a heavily obscured AGN. This relative IR SFR excess (ISX) selection method requires and utilizes redshift information to measure the IR excess effectively, and it appears to isolate a different IR sample from the methods above \\citep{Donley2008,Alexander2008}. A few of the candidates in \\citet{Daddi2007} have been reliably identified as CT AGNs \\citep[][]{Alexander2008}. The ISX sample is also contaminated by star-forming galaxies due to uncertainties in the calculation of SFRs; the heavily obscured AGN fraction in the \\citet{Daddi2007} sample is estimated to be $\\approx25\\%$ \\citep{Alexander2011}, and the rest of the objects are likely star-forming galaxies.  In this paper, we improve the ISX selection method and utilize it to search for \\hbox{X-ray} undetected, heavily obscured AGNs at $z\\approx0.5$--1 in the \\chandra\\ Deep Field-South (\\hbox{CDF-S}; \\citealt{Giacconi2002,Luo2008,Xue2011}) and Extended \\chandra\\ Deep \\hbox{Field-South} (E-CDF-S; \\citealt{Lehmer2005}). We choose to apply the ISX method on the CDF-S and E-CDF-S galaxy sample in this redshift range for the following reasons: \\begin{itemize} \\item {\\it IR-selection of heavily obscured AGNs at $z\\approx0.5\\textrm{--}1$ is a poorly explored territory.} Most of the previous IR selections were focused on samples at $z\\approx2$, including the \\citet{Daddi2007} work. Since AGNs at $z\\approx0.5$--1 contribute significantly to the resolved XRB \\citep[e.g.,][]{Worsley2005,Gilli2007}, and heavily obscured AGNs in this redshift range are expected to play a similar role in the missing fraction of the XRB \\citep[e.g.,][]{Gilli2007,Treister2009}, it is of interest to study the heavily obscured population at these lower redshifts. \\citet{Fiore2009} selected a sample of IR-excess sources in the COSMOS field, which included some objects at $z<1$. However, detailed studies (e.g., the AGN fraction) were not focused on this low-redshift bin. Moreover, the X-ray and IR (24~$\\mu$m) flux limits in the CDF-S field are $\\ga10$ times more sensitive than those in the COSMOS field \\citep{Fiore2009}, allowing identification of the population of heavily obscured AGNs with lower intrinsic luminosities (which are likely more numerous). Recently, \\citet{Georgakakis2010} selected a sample of 19 IR-excess sources at $z\\approx1$ in the AEGIS and GOODS-N fields, but they argued that most of those X-ray undetected candidates are not heavily obscured AGNs based on IR SED modeling.  \\item {\\it The ISX method is best applicable at $z\\approx0.5\\textrm{--}1$.} The ISX method proposed by \\citet{Daddi2007} is a promising technique because of its simple physical motivation and its success in selecting identified CT AGNs at $z\\approx2$. However, the ISX method has not been applied to any other studies, partially due to its known limitations. There are two essential quantities in the ISX method, the IR-based SFR ($\\textrm{SFR}_{\\rm IR+UV}$) and the UV-based SFR ($\\textrm{SFR}_{\\rm UV,corr}$). The IR-based SFR is derived mainly from the observed {\\it Spitzer} MIPS 24~$\\mu$m flux using the \\citet{Chary2001} galaxy SED templates. It has been noted that at $z>1.4$, the IR luminosity (and thus $\\textrm{SFR}_{\\rm IR+UV}$) derived this way is overestimated by an average factor of $\\approx5$ \\citep{Murphy2009,Murphy2011},\\footnote{This is largely due to the fact that the local high-luminosity SED templates in \\citet{Chary2001} do not accurately account for the aromatic features around 24~$\\mu$m and the IR SEDs at $z>1.4$; they show weaker PAH emission than that in high-redshift galaxies by an average factor of $\\approx5$.} leading to additional contamination from galaxies in the ISX sample. The UV-based SFR is calculated using the dust-extinction corrected UV luminosity. \\citet{Daddi2007} used an empirical color-extinction relation to estimate the dust extinction. However, this correlation cannot be applied to relatively old galaxies that are intrinsically red (dominated by old stellar populations instead of being reddened), and all these galaxies were removed from the \\citet{Daddi2007} sample.\\footnote{ \\citet{Daddi2007} removed galaxies with SSFR$<$median(SSFR)/3, where SSFR is the specific SFR, defined as $\\textrm{SSFR}=\\textrm{SFR}_{\\rm IR+UV}/M_{*}$ with $M_{*}$ being the stellar mass. About 15\\% of the sources were excluded this way, the median stellar mass of which is $\\approx5$ times larger than that of the remaining sources.} As AGNs tend to be hosted by massive, red galaxies \\citep[e.g.,][]{Xue2010}, this limitation might significantly affect the completeness of the resulting ISX sample. Both these problems can be substantially alleviated if the ISX method is applied at lower redshifts.\\footnote{We note that applying the ISX method at $z\\approx2$ has the advantage of being able to get a better contrast between star-formation and AGN emission at the observed 24~$\\mu$m wavelength than at lower redshifts.} The 24~$\\mu$m deduced IR luminosity is robust at $z\\la1$ \\citep[e.g.,][]{Elbaz2010}, while the dust extinction can be derived via SED fitting \\citep[e.g.,][]{Brammer2009,Reddy2010}. The \\hbox{SED-fitting} technique is reliable with the high-quality multiwavelength photometric data achievable at $z\\la1$.   \\item {\\it The CDF-S and E-CDF-S are excellent fields in which to perform such a study.} They have been covered by extensive multiwavelength photometric and spectroscopic surveys, including very deep X-ray and 24~$\\mu$m exposures. The COMBO-17 \\citep{Wolf2004,Wolf2008} and Subaru \\citep{Cardamone2010} surveys covering the entire CDF-S and E-CDF-S are particularly valuable for the determination of dust extinction via SED fitting. The 17-filter coverage of COMBO-17 between $3600$~\\AA\\ and $9200$~\\AA\\ and 18 medium-band coverage of Subaru between $4200$~\\AA\\ and $8600$~\\AA\\ are useful for distinguishing between a red dust-free galaxy and a blue dusty galaxy. Reliable photometric redshifts (photo-$z$'s) can also be obtained with the high-quality multiwavelength data available.  \\item {\\it We can assess the possibility of directly detecting the missing population of heavily obscured AGNs with future hard X-ray observatories.} One of the science goals of several future hard X-ray missions, such as the {\\it Nuclear Spectroscopic Telescope Array} ({\\it NuSTAR}; \\citealt{Harrison2010}) and {\\it ASTRO-H} \\citep{Takahashi2010}, is to resolve better the XRB at $\\approx10$--30~keV via detecting heavily obscured and CT AGNs directly in the distant universe. While IR-selected heavily obscured AGN candidates at $z\\approx2$ are likely below their sensitivity limits,\\footnote{Assuming an absorption column density of $N_{\\rm H}=10^{24}$~cm$^{-2}$, the 10--30 keV flux of the intrinsically luminous (\\hbox{2--10} keV luminosity of $2\\times10^{44}$~\\lum) CT AGN HDF-oMD49 at $z=2.21$ \\citep{Alexander2008} will still be below the detection limit of {\\it NuSTAR} with a 1~Ms exposure (the {\\it NuSTAR} sensitivity limit is from \\url{http://www.nustar.caltech.edu/}). The average luminosity of the CT AGN candidates in \\citet{Daddi2007} is about an order of magnitude fainter than that of HDF-oMD49.} candidates at $z\\approx0.5$--1 should be more easily detectable, and we will critically assess such possibilities.  \\end{itemize}  This paper is organized as follows. We describe the multiwavelength data and ISX sample selection in \\S2. In \\S3 we present an X-ray stacking analysis of the ISX sample and estimate the heavily obscured AGN fraction among the sample via simulations. In \\S4, we calculate the space density of the selected heavily obscured AGNs and their contribution to the XRB. We also discuss the feasibility of detecting these ISX sources with future hard X-ray observatories. We summarize in \\S5. Throughout this paper, we adopt the latest cosmology with $H_0=70.4$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.272$, and $\\Omega_{\\Lambda}=0.728$ \\citep{Komatsu2010}. All given magnitudes are in the AB system \\citep[e.g.,][]{Oke1983} unless otherwise stated. ",
        "conclusions": "We have identified a population of heavily obscured/CT AGNs at $z\\approx0.5$--1 in the CDF-S and E-CDF-S utilizing the ISX selection method. The key points from this work are listed below:  \\begin{enumerate} \\item We have improved the ISX method from \\citet{Daddi2007} by deriving the dust extinction via SED fitting and studying sources at lower redshifts with no bias in their IR-luminosity calculation. The ISX method can be well applied to the X-ray undetected galaxies in the CDF-S region at $z\\approx0.5$--1, given the superb multiwavelength data available. The parent sample of 1313 galaxies were defined from an initial sample of 5237 optically and 24~$\\mu$m selected sources in the E-CDF-S, with the requirements of selected redshift interval, excellent multiwavelength coverage, high stellar mass, and no X-ray detection (\\S2).  \\item We have identified 242 ISX sources in the CDF-S and E-CDF-S at $z\\approx0.5$--1; these sources tend to be hosted by evolved galaxies with high stellar mass and little dust (\\S2.2). An X-ray stacking analysis of 23 of the objects in the central CDF-S region resulted in a very hard X-ray signal with an effective photon index of $\\Gamma=0.6_{-0.4}^{+0.3}$, indicating a significant contribution from obscured AGNs (\\S3.1).  \\item We have performed Monte Carlo simulations to estimate the AGN fraction in the ISX sample and assess the intrinsic properties of the obscured AGNs. We modeled the observed X-ray flux considering both the star-formation and AGN contributions, and we utilized the {\\sc MYTorus} model to treat the spectra of heavily obscured AGNs. The requirement that the sources are not individually detected in the 4 Ms CDF-S sets strong constraints on their intrinsic properties. We infer that $74\\pm25\\%$ of the ISX sources are obscured AGNs, within which $\\approx95\\%$ are heavily obscured and $\\approx80\\%$ are CT (\\S3.2).  \\item The heavily obscured (CT) AGNs discovered in our ISX sample have moderate intrinsic X-ray luminosities; the interquartile range of the intrinsic 2--10~keV luminosities is \\hbox{(0.9--4)$\\times10^{42}$}~\\lum\\ [(1--5)$\\times10^{42}$~\\lum]. The space density of the heavily obscured (CT) AGNs is $(2.0\\pm0.7)\\times10^{-4}$~Mpc$^{-3}$ [$(1.6\\pm0.5)\\times10^{-4}$~Mpc$^{-3}$]. These heavily obscured (CT) AGNs are expected to contribute $\\approx 1.4\\%$ ($\\approx1.2\\%$) of the total XRB flux in the 10--30 keV band. In the 6--8 keV band, the 23 ISX sources provide $<1.2\\%$ of the XRB flux. The X-ray undetected optical galaxies in the CDF-S produce a 2.5~$\\sigma$ stacked signal in the 6--8 keV band, accounting for $28\\pm11\\%$ of the XRB flux, which is about the entire unresolved XRB fraction. The space density of the ISX AGNs and their contribution to the XRB could be increased by a factor of $\\approx1.5$--2 due to our conservative ISX definition (\\S4).   \\item These heavily obscured or CT AGNs will probably not be detected by hard X-ray observatories under development such as {\\it NuSTAR} or {\\it ASTRO-H} due to their moderate intrinsic X-ray luminosities and significant obscuration. Most of the hard X-ray sources that will be detected by these facilities are likely already detected in the 4 Ms CDF-S given the extremely high sensitivity of \\chandra\\ to point sources (\\S4.4).   \\end{enumerate}    The ISX selection method presented in this paper can be applied to other survey fields with good optical-to-IR coverage, and it can be expanded to higher redshifts if the IR luminosity can be estimated reliably, e.g., combining MIPS 24~$\\mu$m data with data from deep {\\it Herschel} surveys at 100 and 160~$\\mu$m (e.g., \\citealt{Shao2010}). We will explore these possibilities in future work.  ~\\\\ We acknowledge financial support from CXC grant SP1-12007A (BL, WNB, YQX), CXC grant G09-0134A (BL, WNB, YQX), NASA ADP grant NNX10AC99G (WNB), the Science and Technology Facilities Council (DMA), and ASI/INAF grant I/009/10/0 (AC, CV, RG). We are grateful to T. Yaqoob and K. D. Murphy for providing support for the {\\sc MYTorus} model and making available the data for our desired half-opening angle. We thank A. T. Steffen for helpful discussions. We also thank the referee for carefully reviewing the manuscript and providing helpful comments that improved this work."
    },
    "1107/1107.4242_arXiv.txt": {
        "abstract": "For the first time nonradial oscillations of superfluid nonrotating stars are self-consistently studied at finite stellar temperatures. We apply a realistic equation of state and realistic density dependent model of critical temperature of neutron and proton superfluidity. In particular, we discuss three-layer configurations of a star with no neutron superfluidity at the centre and in the outer region of the core but with superfluid intermediate region. We show, that oscillation spectra contain a set of modes whose frequencies can be very sensitive to temperature variations. Fast temporal evolution of the pulsation spectrum in the course of neutron star cooling is also analysed. ",
        "introduction": "Studying the pulsations of neutron stars (NSs) is very important and actively developing area of research, since comparison of pulsation theory with observations can potentially give valuable information about the properties of superdense matter (\\citealt{afjkkrrz11}). Yet, it is an extremely difficult theoretical problem even for normal (nonsuperfluid) stars. Superfluidity of baryons additionally complicates the theory because in superfluid (SFL) matter one should deal with several independent velocity fields. As a consequence, the hydrodynamics describing pulsations of SFL NSs is much more complicated in comparison to ordinary (normal) hydrodynamics (e.g., \\citealt{BulkVisc}).  In this letter we report on a substantial progress in modeling and understanding the nonradial oscillations of nonrotating SFL NSs in full general relativity. It should be noted that the nonradial oscillations of such stars have been intensively studied in the literature starting from the seminal paper by \\cite{LM94}. In particular, nonradial oscillations of general relativistic NSs were considered by \\cite*{cll99, acl02, yl03, SF_oscil}. Because of the complexity of the problem, most of these papers used a simplified microphysics input, i.e., toy-model equations of state and simplified models of baryon superfluidity (see, however, \\citealt{SF_oscil}; \\citealt*{hap09}; \\citealt{ha10}). Moreover, in {\\it all} these studies SFL matter was treated as being at zero temperature, an assumption that is unjustified at not too low stellar temperatures and may lead to a quantitatively incorrect oscillation spectra (\\citealt{GA, RadPuls}; this letter).  Here we improve on this by considering nonradial oscillations of SFL neutron stars at {\\it finite} temperatures. We follow an approach of \\cite{GK11} (hereafter GK11) which allows us to analyse oscillations of general relativistic NSs employing realistic equation of state, density dependent profiles of nucleon critical temperatures and fully relativistic finite-temperature SFL hydrodynamics.  As it was first found by \\cite{LM94}, oscillation spectrum of a SFL NS consists of two distinct classes of modes, the so called normal and superfluid modes. The frequencies of normal modes almost coincide with the oscillation frequencies of a normal star and hence are independent of temperature. The spectrum of these modes is therefore very well studied in the literature (see, e.g., \\citealt*{ThornPaperI, Osc1, bfg04}). On the contrary, SFL modes can be very temperature-dependent. Using the approach of GK11 they can be decoupled from the normal modes and studied separately. Since radial SFL modes have already been thoroughly analysed in \\cite{RadPuls}, here we focus on the nonradial SFL modes. In what follows, the speed of light $c=1$. ",
        "conclusions": "In this letter we study, for the first time, nonradial oscillations of SFL NSs at finite temperatures. Use of an approach developed in GK11 enable us to solve this problem in full general relativity, employing the realistic equation of state (APR) and realistic, density-dependent profiles of nucleon critical temperatures (Fig.\\ \\ref{Fig:Tcr}). It is shown that equations describing SFL modes can generally be reduced to the simple second-order differential Eq.\\ (\\ref{main}). The eigenfrequency spectrum (Fig.\\ \\ref{Fig:Spec}) and eigenfunctions (Fig.\\ \\ref{Fig:eigen}) for this equation are carefully analysed. It is demonstrated that dependence of eigenfrequencies $\\omega_{ln}$ on $T^\\infty$ is determined by two competing effects: (i) decreasing of $\\omega_{ln}$ with expanding SFL-region and (ii) growing of $\\omega_{ln}$ with increasing of the local speed of SFL sound $v^{\\rm sf}$. These results agree with the conclusions made by Kantor and Gusakov (2011) for a radially oscillating NS.  In addition, we examine the evolution of oscillation spectrum in the course of NS cooling (Fig.\\ \\ref{Fig:l2evol}). We find that acceleration of NS cooling soon after the onset of triplet neutron superfluidity leads to a very fast modification of the spectrum -- the eigenfrequencies can vary dramatically on a time-scale of months.  In the present work we completely ignore various dissipative effects in pulsating SFL NSs. Meanwhile, our results provide a strong basis to study these effects using the perturbative scheme suggested in GK11. In particular, to determine the damping time $\\tau$ of some normal or SFL mode one could proceed in the following steps: (i) find the vectors $u^{\\mu}$ and $w^{\\mu}_{(\\rm n)}$ assuming $s=0$ and neglecting dissipation (see GK11 and Sec.\\ 2); (ii) use them to calculate the dissipative terms entering the SFL hydrodynamics (and hence the rate of change $\\dot{E}$ of the oscillation energy $E$, see, e.g., \\citealt{RadPuls}); (iii) calculate $\\tau$ as $\\tau=-E/(2 \\dot{E})$.  The other important problem concerns gravitational radiation from pulsating SFL NSs. The radiation from normal modes can be accurately calculated already in the $s=0$ limit (and will be the same as that for a nonsuperfluid NS). In turn, SFL modes are not coupled with the metric in the $s=0$ limit, so that to calculate gravitational radiation from them one must use the first-order perturbation theory in $s$. Interestingly, this problem is equivalent to finding gravitational radiation from a {\\it normal} NS experiencing oscillations under an action of a small external force $\\propto s$. An intensity of such radiation will be reduced by a factor of $s^2\\sim 10^{-3}$ in comparison to normal modes with the same $E$. Detailed studies of dissipation in SFL NSs are a very important task, e.g., for calculation of the instability windows of r-modes; we plan to address it in the near future."
    },
    "1107/1107.1717_arXiv.txt": {
        "abstract": "We compare atomic gas, molecular gas, and the recent star formation rate (SFR) inferred from \\ha\\ in the Small Magellanic Cloud (SMC). By using infrared dust emission and local dust-to-gas ratios, we construct a map of molecular gas that is independent of CO emission. This allows us to disentangle conversion factor effects from the impact of metallicity on the formation and star formation efficiency of molecular gas. On scales of 200~pc to 1~kpc (where the distributions of \\htwo\\ and star formation match well) we find a characteristic molecular gas depletion time of $\\tau_{\\rm dep}^{\\rm mol} \\sim 1.6$~Gyr, similar to that observed in the molecule-rich parts of large spiral galaxies on similar spatial scales. This depletion time shortens on much larger scales to $\\sim 0.6$~Gyr because of the presence of a diffuse \\ha\\ component, and lengthens on much smaller scales to $\\sim 7.5$~Gyr because the \\ha\\ and \\htwo\\ distributions differ in detail. We estimate the systematic uncertainties in our dust-based $\\tau_{\\rm dep}^{\\rm mol}$ measurement to be a factor of $\\sim 2$--$3$. We suggest that the impact of metallicity on the physics of star formation in molecular gas has at most this magnitude, rather than the factor of $\\sim 40$ suggested by the ratio of SFR to CO emission. The relation between SFR and neutral ($\\htwo+\\hi$) gas surface density is steep, with a power-law index $\\approx2.2\\pm0.1$, similar to that observed in the outer disks of large spiral galaxies.  At a fixed total gas surface density the SMC has a $5-10$ times lower molecular gas fraction (and star formation rate) than large spiral galaxies. We explore the ability of the recent models by \\citet{KRUMHOLZ09} and \\citet{OSTRIKER10} to reproduce our observations. We find that to explain our data at all spatial scales requires a low fraction of cold, gravitationally-bound gas in the SMC. We explore a combined model that incorporates both large scale thermal and dynamical equilibrium and cloud-scale photodissociation region structure and find that it reproduces our data well, as well as predicting a fraction of cold atomic gas very similar to that observed in the SMC. ",
        "introduction": "The relation between gas content and star formation activity in galaxies has been a matter of intense investigation since the pioneering work of Maarten Schmidt. \\citet{SCHMIDT59} suggested that the star formation rate (SFR) in a galaxy is proportional to a power of the gas density, such that $\\rho_{\\rm SFR}\\sim\\rho_{\\rm gas}^n$, where $n\\simeq1-3$ and most likely $n=2$ based on a number of arguments that included the luminosities of open clusters, the abundance of helium, and the vertical distribution of objects in the plane of the Milky Way.  Here we refer to the quantitative relation between gas and star formation as the ``star formation law'' for convenience without intending to suggest a physical law or a specific functional form.  Most modern empirical studies of the extragalactic star formation law follow those by \\citet{KENNICUTT89,KENNICUTT98}, who studied disk-averaged correlations in a sample of $\\sim100$ galaxies including starbursts and high-mass dwarfs. This work linked the surface density of star formation rate, $\\Ssfr$, to the surface density of total neutral (atomic and molecular) gas, $\\Sgas=\\Shi+\\Shtwo$. \\citet{KENNICUTT98} found that $\\Ssfr\\propto\\Sgas^{1+p}$ with $1+p\\approx {1.4\\pm0.15}$ for his composite sample of galaxies. Since the gas depletion time $\\tau_{\\rm dep}^{\\rm gas}\\equiv \\Sgas/\\Ssfr \\propto \\Sgas^{-p}$ for a power-law relation, $p=0$ indicates that the gas depletion time is constant (independent of environment), while $p>0$ indicates that star formation is more rapid in high-$\\Sgas$ regions. In practice, it is also important to keep in mind the observational complexities associated with measuring \\Sgas\\ and particularly \\Shtwo\\ from CO observations, and the systematics thus introduced.  Several subsequent studies focused on the star formation law within galaxies, employing high resolution molecular and \\hi\\ observations \\citep{MARTIN01,WONG02,BOISSIER03}. More recently, \\citet{BIGIEL08} and \\citet{LEROY08} analyzed 12 nearby spirals with high quality \\hi, CO, far-infrared, and far-ultraviolet data. These observations show a clear correlation between $\\Ssfr$ and the surface density of molecular gas, $\\Shtwo$, with an approximately constant star formation rate per unit molecular gas, yielding an approximately linear power-law index $1+p=1.0\\pm0.2$. By contrast, they found a steep correlation between $\\Ssfr$ and the surface density of atomic gas, $\\Shi$ \\citep[the observed correlation in the optical disks of large galaxies has $1+p\\gtrsim 2.7-3.5$, and $1+p\\sim2$ in the \\hi-dominated outer disks][]{BIGIEL08,BIGIEL10b}. \\citet{SCHRUBA11} and \\citet{BIGIEL11} extended these findings to a larger sample of 30 disk galaxies, while \\citet{BLANC09} and \\citet{RAHMAN11} arrived at similar conclusions in very detailed studies of individual targets.  A consequence of the observed linearity of the star formation law over the molecule-dominated regions of disks is that the typical time scale to deplete the molecular gas by star formation in a disk galaxy is $\\tau_{\\rm dep}^{\\rm mol}=\\Smol/\\Ssfr\\sim1.9\\pm0.9$ Gyr (\\Smol\\ corresponds to \\Shtwo\\ corrected by a 1.36 factor to account for the cosmic abundance of helium).  The lack of dependence of $\\tau_{\\rm dep}^{\\rm mol}$ on environment can be naturally understood if two conditions are fulfilled.  The first condition is for star formation to be a local process primarily determined by the conditions inside the giant molecular clouds (GMCs) where it takes place. GMCs are themselves isolated from the global galactic environment provided that they are self-gravitating and therefore overpressured with respect to the neighboring gas \\citep{MO07}. The second condition is for the properties of GMCs to be universal, and therefore independent of their galactic environment. Because star formation takes place in the densest regions of GMCs, themselves self-gravitating and thus mostly decoupled from their surroundings, the first condition appears likely \\citep{KRUMHOLZ05}, at least in mid-disk and outer-galaxy regions.  In galactic centers it is less clear whether there are isolated, gravitationally bound GMCs, or instead a distributed molecular medium. Environmental considerations may be more important if cloud--cloud collisions play an important role \\citep{TAN00}. Similarly, the second condition appears to be broadly satisfied as the resolved properties of GMCs (or at least of the dense regions of GMCs, bright in CO) are observed to be remarkably constant over a wide range of galaxy environments (the Local Group spirals and nearby dwarfs, Bolatto et al. 2008\\nocite{BOLATTO08}; the Milky Way, Heyer et al. 2009\\nocite{HEYER09}; the LMC, Hughes et al. 2010\\nocite{HUGHES10}; the outer disk of M33, Bigiel et al. 2010\\nocite{BIGIEL10a}).  Most of the studies of the relation between gas and star formation, and the ensuing conclusions, focus on large galaxies which tend to be rich in molecules. There is a dearth of comparable information for low-mass low-metallicity star-forming dwarf galaxies, primarily because such objects emit only very faintly in molecular gas tracers such as CO. The measurements that do exist reveal large ratios of SFR to CO emission in late-type, low-mass galaxies \\citep[e.g.,][]{YOUNG96,GARDAN07,LEROY07B,KRUMHOLZ11}.  Studies of the star formation law in dwarf galaxies are interesting for a number of reasons. Fundamentally they probe a very different physical regime from large spiral galaxies, one where atomic gas dominates the surface density of the interstellar medium (ISM) on large scales, and elements heavier than helium are less abundant. This dearth of heavy elements causes a number of differences in gas chemistry and physical conditions. For example, dust-to-gas ratios are lower, producing lower extinctions and higher photodissociation rates \\citep[the likely cause of their general CO faintness,][]{ISRAEL86,MALONEY88,LEQUEUX94,BOLATTO99}.  More importantly, studies of the star formation law in low-mass low-metallicity star-forming dwarf galaxies provide fertile testing ground for star formation theories.  A natural implication of the approximately constant molecular $\\tau_{\\rm dep}^{\\rm mol}$ among normal galaxies is that the effectiveness at forming molecular gas plays a critical role in establishing the star formation rate. The molecular gas fraction varies systematically both within and among galaxies \\citep{YOUNG91,WONG02,BLITZ06,LEROY08}. Dwarf galaxies appear to have low molecular fractions, despite their large gas fractions and long time scales to deplete gas reservoirs at current rates of star formation. If their $\\tau_{\\rm dep}^{\\rm mol}$ values are similar to normal galaxies, it would suggest that GMC formation (and destruction) is the rate-limiting step for star formation in dwarf systems.  Because their low metallicities contrast with those in spiral galaxies, star-forming dwarf galaxies can break a number of degeneracies in the physical drivers of the molecular fraction. Different models predict very different behaviors for the impact of metallicity on molecular fractions and star formation rates according to their emphasis on dynamics, thermodynamic equilibrium, or dust shielding.  For example, the requirement of a minimum extinction for star formation to occur \\citep[e.g.,][]{MCKEE89,LADA09} would have a proportionally more important impact in low metallicity objects, while models that are solely based on dynamics cannot distinguish between low and high metallicity.  The main obstacle to using dwarf galaxies to test various aspects of the star formation law is the difficulty in observing their molecular gas distribution, stemming from the faintness of their CO emission and the uncertainty in its quantitative relation to \\htwo. The Small Magellanic Cloud (SMC) offers a unique opportunity in this regard. Its proximity allows us to probe small spatial scales with observations of modest angular resolution. As a result, we can use dust emission to construct a map of $\\Sigma_{\\rm H2}$ that does not depend on CO. In this study we use such a map to compare the distributions of molecular gas, atomic gas, and recent star formation traced by ionizing photon production.  \\subsection{The Small Magellanic Cloud as a Galaxy} \\label{subsec:smcintro}  By virtue of its location and properties the SMC is a prime laboratory for the study of the relation between gas and star formation in low-mass galaxies. Indeed, the SMC was the target of the first extragalactic study of the star formation law, relating \\hi\\ and stellar surface densities \\citep{SANDULEAK69}.  Located scarcely 61~kpc away \\citep{HILDITCH05,KELLER06,SZEWCZYK09}, the SMC is the nearest gas-rich low metallicity dwarf galaxy with active star-formation \\citep[$Z_{\\rm SMC}\\sim Z_\\odot/5$,][]{DUFOUR84,KURT99,PAGEL03}. As such, it provides invaluable insight into the physics and chemistry of the ISM in chemically primitive star-forming galaxies. Moreover, because of its proximity it is possible to carry out studies on the scale of individual young stellar objects, which reveal subtle differences in the temperature and chemistry of star-forming cores \\citep{VANLOON10b,OLIVEIRA11}.  With an \\hi\\ mass of $M_{\\rm HI}\\simeq4.2\\times10^8$~\\msun\\ \\citep{STANIMIROVIC99} the SMC is rich in atomic gas, although also strikingly defficient in cold \\hi\\ \\citep{DICKEY00}. It was already apparent in the early observations that the distribution of the \\hi\\ is complex, with more than one component along the line-of-sight \\citep{MCGEE81}. This complexity is at least in part due to the presence of several supergiant shells \\citep{STANIMIROVIC99}. The underlying dynamics appear to be disk-like with an inclination $i\\approx40^\\circ\\pm20^\\circ$, although disturbed by the interaction with the Milky Way and the LMC \\citep{STANIMIROVIC04}. In our analysis we will ignore this complexity, but the geometry of the SMC and the level of turbulent support in the gas remain some of the most important caveats in the comparison to star formation models.  The SMC hosts a healthy amount of star formation despite being disproportionately faint in CO emission \\citep{KENNICUTT95,ISRAEL86,ISRAEL93}.  In fact, in the SMC the CO is under-luminous with respect to the star formation activity by almost two orders of magnitude when compared to normal disks, or even more massive small galaxies.  A possible explanation for this fact is that the SMC is extraordinarily efficient at turning the available molecular gas into stars (i.e., the H$_2$ depletion time is very short), which would suggest that we are either observing an out-of-equilibrium situation (a fleeting starburst) or that the conditions conspire to keep a small reservoir of extremely short-lived GMCs.  A more likely alternative is that the weak CO emission is not representative of star-forming molecular gas, with H$_2$/CO considerably larger than in the Milky Way \\citep{ISRAEL86,RUBIO93b}.  \\citet{LEROY07} used new far-infrared observations to map the H$_2$ distribution in the SMC bypassing CO emission, following an extension of the methodology previously employed by \\citet{ISRAEL97} in the Magellanic Clouds. The feasibility of using dust to map \\htwo\\ was clearly demonstrated by \\citet{DAME01} in the Milky Way \\citep[see also][for an early study of the correlation between dust emission and gas in the Milky Way]{BLOEMEN90}. It was also shown to be consistent with virial masses on large scales \\citep{LEROY09}. Furthermore, a systematic application of an extension of this methodology throughout the Local Group produces molecular masses that are consistent with those obtained by other methods at approximately Galactic metallicities \\citep{LEROY11}.  The analysis by \\citet{LEROY07} shows that the column density of molecular gas present in the SMC is far in excess of that derived by applying the Galactic CO-to-H$_2$ conversion factor, \\xco, to the CO observations.  Their modeling yields a total molecular mass $M_{\\rm H2}\\approx3.2\\times10^7$ M$_\\odot$. Therefore the neutral ISM in the SMC is approximately 10\\% molecular, a lower fraction than observed in normal late spirals but not dramatically so \\citep{YOUNG91}. Further refinements to the analysis \\citep[][and this work]{LEROY09,LEROY11} broadly confirm these numbers.  In this study we analyze the spatially resolved correlation between the distributions of molecular gas, atomic gas, and recent star formation traced by ionizing photon production on different scales, comparing the star formation law in this small low-metallicity galaxy with that in large disks. We present our data and discuss the methodology we use to measure molecular column densities and star formation rates in \\S\\ref{sec:method}. We show our main observational results in \\S\\ref{sec:results}, focusing on the relation between \\htwo\\ and star formation in \\S\\ref{subsec:molsfr}, total gas and star formation in \\S\\ref{subsec:gassfr}. In \\S\\ref{sec:comparison}, we compare our results to recent analytical physical models of star formation in galaxies. Finally, we summarize our conclusions in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We present a detailed analysis of the correlation between the atomic and molecular gas phases in the SMC, and the star formation activity. This is the first time such analysis has been done at high spatial resolution in a galaxy of such a low metallicity. We use modeling of the {\\em Spitzer} dust continuum observations \\citep{BOLATTO07,GORDON11} to calculate the \\htwo\\ surface densities on scales of $\\theta\\sim40\\arcsec$ or $r\\sim12$~pc, together with observations of the \\hi\\ \\citep{STANIMIROVIC99} and \\ha\\ \\citep{SMITH99} distributions (Figure \\ref{fig:h2map}). An important caveat to keep in mind when considering the interpretations of the results is the complexity of the \\hi\\ distribution in the SMC, referred to in \\S\\ref{subsec:smcintro}. Another important caveat are the uncertainties involved in the production of the \\htwo\\ map, which we discuss in \\S\\ref{subsec:H2map}.  We find that, in the regions where we detect \\Shtwo, we measure a typical molecular gas depletion time $\\tau_{\\rm dep}^{\\rm mol}\\sim7.5$ Gyr on scales $r\\sim12$~pc with a factor of 3.5 uncertainty accounting for the scatter as well as the systematics associated with our \\htwo\\ map as well as the geometry of the source (Figure \\ref{fig:sflmol}). The depletion time shortens when measured on larger spatial scales. On scales $r\\sim200$~pc (a typical scale over which molecular gas and \\ha\\ emission are well correlated in galaxies) and $r\\sim1$~kpc (a typical scale for extragalactic studies) we measure $\\tau_{\\rm dep}^{\\rm mol}\\approx1.6$~Gyr. The molecular depletion time for the SMC as a whole is shorter, $\\tau_{\\rm dep}^{\\rm mol}\\sim0.6$~Gyr, due to a large component of extended low level \\ha\\ emission.  These results are consistent with the typical depletion time of $\\tau_{\\rm dep}^{\\rm mol}\\sim2$ Gyr observed in normal disks on kpc scales \\citep{BIGIEL08,BIGIEL11}. Consequently the relation between molecular gas and star formation activity appears to be at most only weakly dependent on metallicity. This finding suggests that the molecular content can be used to infer the star formation activity (and viceversa) even in galaxies that are chemically primitive and deficient in heavy elements.  We also measure the relation between star formation and total gas surface density, which is dominated by \\hi\\ over most of the SMC (Figure \\ref{fig:sflgas}). We find that the relation is similarly steep in the SMC ($1+p=2.2\\pm0.1$) and in the outer disks of normal galaxies \\citep[$1+p\\approx2$; ][]{BIGIEL10b}, but it is displaced toward much larger surface densities. In the SMC the \\hi\\ surface density reaches values as high as $\\Shi\\sim100$ \\msunperpcsq, while in most galaxies $\\Shi\\lesssim10$ \\msunperpcsq. As a consequence the use of the standard total gas to star formation rate relation dramatically overpredicts the star formation activity over much of this source. This finding supports the explanation that DLAs are deficient in star formation because of very low molecular fractions \\citep{WOLFE06}. The gas (or \\hi) depletion time for the SMC is approximately a Hubble time, $\\tau_{\\rm dep}^{\\rm gas}\\approx11.8$~Gyr.  We compare our results with the recent analytical models by \\citet[][KMT09]{KRUMHOLZ09} and \\citet[][OML10]{OSTRIKER10}. For high-resolution observations, we find that KMT09 is very successful at reproducing the displacement in the relation between \\Sgas\\ and \\Ssfr\\ toward higher surface densities (Figure \\ref{fig:fig3}). We also find that with a clumping factor $c$ of unity it approximately predicts the observed relation between the molecular to atomic ratio $\\Shtwo/\\Shi$ and \\Sgas\\ (Figure \\ref{fig:fig5}), although on a linear scale it becomes apparent that there is a systematic overprediction at high \\Sgas\\ by $0.3-0.5$ dex (Figure \\ref{fig:kmtcomp}).  The slope of \\Ssfr\\ vs. \\Sgas\\ predicted by KMT09 for fixed $cZ$ is steeper than the relation observed in the SMC, both at the observed 12 pc resolution and when the data is averaged over 200 pc and 1 kpc scales (Figures \\ref{fig:fig3} and Figure \\ref{fig:fig6}).  Given that KMT09 co-locates \\hi\\ and \\htwo\\ within the same dense cloud complexes, a strong relation between \\Shi\\ and \\Ssfr\\ would be expected at 200 pc scales, but this is not seen in the observations (Figure \\ref{fig:fig4}).  We test the effect of smoothing the observations to larger spatial scales, and find that changing the clumping parameter in the KMT09 model does not match the effect of smoothing the observations (Figure \\ref{fig:fig6}). The bulk of the total gas surface densities in the SMC stay approximately constant for spatial scales in the $r\\sim12$~pc---1~kpc range.  This suggests that much of the \\hi\\ is in a warm, diffuse component, which stands in contrast to the starting assumption of KMT09 that the \\hi\\ consists primarily of cold gas concentrated in the envelopes of dense atomic-molecular complexes, with the warm \\hi\\ a small fraction of the total ISM mass.  In comparing the observations to the predictions of OML10, we find that the expected \\Sgbc/\\Sdiff\\ is very large in comparison to the observed \\Shtwo/\\Shi\\ in the SMC (Figure \\ref{fig:fig5}), if we adopt the same heating efficiency and turbulence parameters for thermal/dynamical equilibrium as in typical large spirals with $Z\\sim 1$. In other words, the model conspicuosly overpredicts the molecular fraction in the SMC for the default parameters. This motivates us to investigate variations and extensions of the OML10 model.  In the first variation, which we call OML10h, we introduce a metallicity dependence in the ratio of gas heating to star formation that produces enhanced heating at low metallicities.  This modification is based on the idea that a lower dust-to-gas ratio allows UV to escape and travel farther from star-forming regions, potentially raising $G_0'$ significantly in the diffuse phase.  An enhanced $G_0'$ leads to increased heating in the diffuse medium, permitting thermal equilibrium of warm \\hi\\ at high pressure.  Without enhanced heating, warm gas at the typical thermal pressure in the SMC ($P_{\\rm th}/k\\sim 3\\times 10^4$~K~cm$^{-3}$ for \\Shi$\\sim 65$ \\msunperpcsq) would cool very quickly.  Enhanced heating changes the balance between the diffuse and self-gravitating phases of the gas.  As a consequence the self-gravitating phase is less abundant than at normal metallicities under similar conditions for \\Sgas. This model does a very reasonable job at matching the $\\Shtwo/\\Shi$ derived for the SMC (Figures \\ref{fig:fig5}), although there are still local discrepancies at the level of $\\sim\\pm0.3$ dex.   As an extension of OML10, which we term OML10p, we do not change the balance between the diffuse and self-gravitating components, but add an accounting for the abundance of \\hi\\ within gravitationally bound clouds, based on the photodissociation formalism of KMT09.  The \\hi\\ is split between the diffuse and the self-gravitating phases, and under the typical SMC conditions, half of the \\hi\\ would be self-gravitating.  The photodissociation calculation for the GBC component can also be applied to the enhanced-heating model; we denote this as model OML10ph.  Model OML10ph shares the reduced abundance of the self-gravitating phase with model OML10h, and includes a contribution to \\hi\\ that originates in cold self-gravitating clouds. Note that in adding the photodissociation estimate for GBCs, we do not change the predicted SFR.  Thus, model OML10 and OML10p have the same \\Ssfr\\ vs. \\Sgas\\ as each other, as do model OML10h and OML10ph (Figure \\ref{fig:fig11}).  When compared to the observed \\Shtwo/\\Shi\\ in the SMC, model OML10p slightly overpredicts the molecular ratio ($\\sim 10-30\\%$ for the model vs. $\\sim 5-15\\%$ for the data; see Figure \\ref{fig:fig10}).  For the typical parameters of the SMC, $\\sim 30-60\\%$ of the total \\hi\\ would be in GBCs for model OML10p (Figure \\ref{fig:fig9}).  Although this is less than in KMT09 (which puts all of the \\hi\\ in the envelopes of cold, dense cloud complexes), to some extent it shares the difficulty that in observations \\hi\\ is not correlated with star formation activity (Figure \\ref{fig:fig4}).  Model OML10ph, with a lower abundance of self-gravitating gas and hence \\Shtwo\\ due to enhanced heating, follows the observed \\Shtwo/\\Shi\\ magnitude and slope better than model OML10p (Figure \\ref{fig:fig10}).  OML10ph also has a lower ratio of cold to warm \\hi\\ than OML10p or KMT09, consistent with the observations of \\citet{DICKEY00}.  When compared to star formation in the SMC, the enhanced-heating models (OML10h, OML10ph) fit the relation between \\Ssfr\\ and \\Sgas\\ much better than the models that adopt the same heating efficiency as higher-metallicity spirals (OML10, OML10p), as shown in Figure \\ref{fig:fig11}.  In essence, the observed star formation rate requires that the eligible cold, self-gravitating gas (whether \\hi\\ or \\htwo) is less abundant at lower metallicities, otherwise we would expect larger star formation rates for the observed \\Sgas\\ surface densities.  The diminished abundance of the cold self-gravitating phase is supported by observations of the distribution of \\hi\\ temperatures \\citep{DICKEY00}, and by the relative uniformity of \\Shi\\ when averaged at different scales.  The possibility of an enhanced radiation field that could maintain warm, diffuse \\hi\\ at high surface density and pressure is suggested by recent dust emission modeling \\citep{SANDSTROM10}. Further observational confirmation of enhanced heating could come from observations of the cooling in the diffuse atomic gas. Further characterization of the fraction of cold atomic gas and its spatial distribution in both Magellanic Clouds using \\hi\\ absorption/emission techniques would also be extremely valuable in constraining the models."
    },
    "1107/1107.4583_arXiv.txt": {
        "abstract": "Ibata et al. reported evidence for density and kinematic cusps in the Galactic globular cluster M54, possibly due to the presence of a $9400~M_\\odot$ black hole.  Radiative signatures of accretion onto M54's candidate intermediate-mass black hole (IMBH) could bolster the case for its existence.  Analysis of new {\\em Chandra\\/} and recent {\\em Hubble Space Telescope\\/} astrometry rules out the X-ray counterpart to the candidate IMBH suggested by Ibata et al.  If an IMBH exists in M54, then it has an Eddington ratio of $L(0.3-8~{\\rm keV}) / L({\\rm Edd}) < 1.4 \\times 10^{-10}$, more similar to that of the candidate IMBH in M15 than that in G1.  From new imaging with the NRAO Very Large Array, the luminosity of the candidate IMBH is $L(8.5~{\\rm GHz}) < 3.6 \\times 10^{29}$~ergs~s$^{-1}$ (3 $\\sigma$). Two background active galaxies discovered toward M54 could serve as probes of its intracluster medium. ",
        "introduction": "\\label{motivation}  Quantifying the space density of low-mass black holes (BHs) in the local Universe has important implications for predicting gravity wave signals \\citep{hug09}, for understanding formation channels for seed BHs \\citep{vol08}, and for testing simulations of gravitational wave recoil \\citep{hol08}.  Moreover, establishing the existence or break-down of scaling relations between central BHs and galaxies at the low-mass end will help identify the physical processes driving the relation between BH mass $M_{\\rm BH}$ and stellar velocity dispersion $\\sigma_\\ast$ \\citep[e.g.,][]{gre10}.  Planned NIR telescopes in the 30 m class will eventually offer access to BHs with $M_{\\rm BH} \\gtrsim 10^5~M_\\odot$ in the local Universe \\citep{fer03,fer05}.  But lower-mass BHs will remain difficult to find, as their observational signature on surrounding stars is limited to a prohibitively small spatial scale beyond the Local Group.  We refer to these systems as intermediate-mass black holes (IMBHs), since they occupy the gap in mass between the well-studied stellar-mass BHs with $M_{\\rm BH} \\sim 10~M_\\odot$ and the well-established BHs with $M_{\\rm BH} \\gtrsim 10^5~M_\\odot$.  BH masses of order $10^4~M_\\odot$ have been estimated for two globular clusters, namely G1 in M31 \\citep{geb02} and $\\omega$~Cen in the Galaxy \\citep{noy10}.  Both estimates have been controversial.  For G1 the objections of \\citet{bau03} were countered by \\citet{geb05}.  For $\\omega$~Cen \\citet{van10} provide a contradicting mass upper limit.  \\citet{iba09} reported the detection of stellar density and kinematic cusps in M54 (NGC\\,6715), a globular cluster at a distance of 26.3 kpc in the center of the Sagittarius dwarf galaxy \\citep{bel08}.\\footnote{1\\arcsec\\, = 0.13 pc. The core, half-mass and limiting radii of M54 are 6.6\\arcsec\\, = 0.86 pc, 29.4\\arcsec\\, = 3.8 pc and 6.3\\arcmin\\, = 49 pc, respectively \\citep{har96,iba09,sol10}.}  Ibata et al. interpreted their findings as possible evidence for a central BH with $M_{\\rm BH} \\sim 9400~M_\\odot$, similar to G1 and possibly similar to $\\omega$~Cen. But Ibata et al. also noted that the orbits of the cusp stars in M54 could have moderate radial anisotropies.  They demonstrated that the strong density gradient in a cusp could impede isotropic orbits from being established, contrary to the common view that isotropic orbits always arise when cluster relaxation times are relatively short. If such anisotropic orbits are present, they could mimic the signature of a central IMBH.  The interpretation of M54's cusp data is thus ambiguous, and other lines of evidence for, or against, an IMBH should be sought. Simulations suggest a globular cluster with primordial binary stars can achieve a ratio of core to half-mass radius of only 0.08 or less, whereas larger ratios can be realized by adding an IMBH \\citep{tre07,fre07}.  For M54 this ratio is about 0.2 \\citep{har96}. Also, simulations of massive disk galaxies suggest that several tidally-stripped cores of dwarf satellites could retain their IMBHs \\citep{bel10}.  This fits in with M54's location in the center of the disrupting Sagittarius dwarf galaxy \\citep{bel08}.  Radiative signatures of accretion onto M54's candidate IMBH could bolster the case for its existence.  \\citet{ram06a} detected several {\\em Chandra\\/} sources within M54's half-mass radius and suggested, based on the sources' X-ray luminosities and colors, that they were cataclysmic variables or low-mass X-ray binaries.  However, they noted that one possibly-blended source, X-ray Id 2, was located within 1\\arcsec\\, (0.13 pc) of the center of the stellar density, taken a proxy for the center of mass.  This led \\citet{iba09} to suggest that X-ray Id 2 could be the counterpart to the candidate IMBH in M54.  In this paper, we compare recent subarcsecond astrometry for the stellar density center \\citep{gol10} to improved astrometry for X-ray Id 2 \\citep{eva10}, ruling out an X-ray counterpart to the candidate IMBH.  We also present new photometry of M54 with the NRAO Very Large Array (VLA)\\footnote{Operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc.}  \\citep{tho80}, detecting neither the candidate IMBH nor any stellar emitters within the cluster's half-mass radius.  Two X-ray sources with radio counterparts are detected within M54's limiting radiius.  The new X-ray and radio data are presented in \\S~\\ref{data} and their implications are explored in \\S~\\ref{implications}.  A summary and conclusions appear in \\S~\\ref{sumcon}. ",
        "conclusions": "\\label{sumcon}  Radiative signatures of accretion onto M54's candidate IMBH, with mass $\\sim 9400~M_\\odot$ \\citep{iba09}, could strengthen the case for its existence.  Comparison of new X-ray and recent optical astrometry does not support the suggestion by \\citet{iba09} that X-ray Id 2 is the counterpart to the candidate IMBH.  Rather, the Eddington ratio of the candidate IMBH is constrained to be less than $1.4 \\times 10^{-10}$, implying a very different accretion state than that of G1's candidate IMBH.  The 8.5-GHz luminosity of M54's candidate IMBH is less than $3.6 \\times 10^{29}$~ergs~s$^{-1}$.  Two background active galaxies behind M54 could serve as probes of the fuel available to the candidate IMBH."
    },
    "1107/1107.5120_arXiv.txt": {
        "abstract": "\\baselineskip 11pt We model the chameleon effect on cosmological statistics for the modified gravity $f(R)$ model of cosmic acceleration.   The chameleon effect, required to make the model compatible with local tests of gravity, reduces force enhancement as a function of the depth of the gravitational potential wells of collapsed structure and so is readily incorporated into a halo model by including parameters for the chameleon mass threshold and rapidity of transition.   We show that the abundance of halos around the chameleon mass threshold is enhanced by both the merging from below and the lack of merging to larger masses. This property also controls the power spectrum in the nonlinear regime and we provide a description of the transition to the linear regime that is valid for a wide range of $f(R)$ models. ",
        "introduction": "\\label{sec:intro}  The modified action $f(R)$ model for cosmic acceleration provides a concrete framework under which to relate local and cosmological tests of gravity. Here the Einstein-Hilbert action is augmented with a general function of the Ricci scalar curvature in such a way as to mimic the cosmological constant at cosmologically low curvature \\cite{Capozziello:2002rd,NojOdi03,Caretal03}.  The extra scalar degree of freedom $d f/dR$ mediates an enhanced gravitational force on scales smaller than its Compton wavelength. In order to hide this enhancement from local tests of gravity, viable models employ the chameleon mechanism \\cite{Mota:2003tc,khoury04a} where the Compton wavelength can shrink in regions of deep gravitational potential wells.   If one assumes that the $f(R)$ model is a complete description of gravity from solar system scales to cosmology, local tests in conjunction with minimal assumptions about the galactic potential place a stringent bound on $f(R)$ models.   Namely the cosmological amplitude of the field $|f_{R0}|$ must be less than a few times $10^{-6}$ \\cite{HuSaw07a}.  Currently the tightest cosmological bounds come from the lack of an excess in the abundance of the most massive dark matter halos or galaxy clusters \\cite{Schmidt:2009am,Lombriser:2010mp}.  The bound on the field amplitude depends on the model but generally lies above the $10^{-5}$ level for models that seek to diminish the modification at high redshift \\cite{Ferraro:2010gh}. Thus cosmological constraints currently are at best approaching the level of solar system tests.  On the other hand they test these modifications on a vastly different scale and it  is possible that $f(R)$ is just an effective theory only strictly valid on cosmological scales.  As cosmological data continue to improve, they will soon begin to probe a regime where cosmology is directly competitive with solar system tests in the full $f(R)$ context. Here the chameleon mechanism becomes important for all halos of mass comparable to or greater than the Galaxy.    This regime is substantially more difficult to characterize due to the non-linearity in the field equation for the enhanced force.  Under the halo model approach, the first step in understanding cosmological statistics associated with the observable properties of dark matter halos is to characterize the mass function. Previous attempts to model the mass function of simulations have not been able to capture the abundance of intermediate mass halos once the chameleon mechanism becomes active \\cite{Halopaper}.    This abundance is doubly enhanced since the extra force augments merging of  low mass halos into intermediate mass halos whereas the chameleon effect shuts down merging of intermediate mass halos to high mass halos.  Based on a mass function model, we can build a parameterized post-Friedmann (PPF) description of the chameleon effect in cosmological statistics such as the dark matter power spectrum. Whereas previous PPF approaches have been based on density thresholds  \\cite{HuSaw07b} and are appropriate for models that utilize the Vainshtein mechanism, they require extensive ad hoc modifications for the chameleon mechanism \\cite{Koyama:2009me}. The problem is that the chameleon mechanism should be parameterized in terms of a proxy for the depth of  gravitational potential wells, which in the halo model can be readily associated with the mass of dark matter halos.     The outline of the paper is as follows.  In \\S \\ref{sec:frsims}, we briefly review the $f(R)$ model and simulations.  We parameterize the chameleon effect in the mass function in \\S \\ref{sec:massfunction}, which we use in \\S \\ref{sec:powerspectrum} to characterize the nonlinear dark matter power spectrum.   We discuss these results in \\S \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In the $f(R)$ model, the chameleon mechanism mediates a restoration of general relativity and ordinary Newtonian forces in deep gravitational potential wells. We have shown that the main impact of the chameleon mechanism on cosmological statistics that depend on the dark matter halo distribution can be simply described by a universal scaling for the transition between modified and unmodified forces in mass or gravitational potential. This type of transition should be contrasted with phenomenological approaches that implement the transition as a function of physical scale (e.g.~\\cite{Amendola:2007rr}) or with models that implement the Vainshtein mechanism where the transition is a function of density \\cite{HuSaw07b}.  In the mass function, the chameleon mechanism leads to a doubly enhanced abundance near the transition.  In the simulations, this enhancement is associated with smaller mass halos still growing to the transition mass due to enhanced forces but transition mass halos no longer merging into high mass halos due to the restoration of ordinary forces.   In our description, this mass conservation property is enforced through the Press-Schechter assumption that all of the dark matter is in halos of some mass.   With this assumption, we can fit the mass function results across a wide range in $f(R)$ models with two constants: the scaling of the transition mass to the background field and the rapidity of the transition in mass.  With a calibration of the mass function and the simulation result that halo profiles as a function of mass are largely unchanged, we can construct predictions for the $N$-point functions.   In particular for the power spectrum, including the chameleon effects on the mass function brings predictions for the power spectrum excess down in agreement with the simulations.   We have provided a simple modification to the Halofit prescription to bridge the linear and nonlinear regimes.     For the higher $N$-point functions, the model predicts that the results  in the deeply nonlinear regime should scale mainly with the single degree of freedom of the mass function rather than carry a unique signature of $f(R)$ gravity \\cite{GSHJV11}.  Moreover, by describing these effects as enhancements over the $\\Lambda$CDM mass function and power spectrum with a physically well-motivated extrapolation to low masses and small scales rather than an absolute prediction of $f(R)$ statistics, one can use state-of-the-art simulations and mass calibrations for $\\Lambda$CDM predictions and look for a parameterized excess over those in the data (see, e.g.~\\cite{Schmidt:2009am}).   Uncertainties in our $f(R)$ prescription merely translate into uncertainties in the $f(R)$ parameters rather than a false positive detection of modified gravity.    These techniques provide simple but approximate means for testing the $f(R)$ model with weak lensing and other statistics that require large dynamic range and precision.  Should in the future a positive detection occur, then high resolution gas simulations can be performed to refine the calibration of cosmological observables in the $f(R)$ model.       \\smallskip \\noindent {\\it Acknowledgments}:   We thank Simone Ferraro, Surhud More and Fabian Schmidt for useful discussions. YL and WH were supported by the Kavli Institute for Cosmological Physics (KICP) at the University of Chicago through grants NSF PHY-0114422 and NSF PHY-0551142 and an endowment from the Kavli Foundation and its founder Fred Kavli.   WH was additionally supported by U.S.~Dept.\\ of Energy contract DE-FG02-90ER-40560 and the David and Lucile Packard Foundation. Computational resources for the cosmological simulations were provided by the KICP-Fermilab computer cluster.    \\vfill"
    },
    "1107/1107.3917_arXiv.txt": {
        "abstract": "For very general scalar-field theories in which the equations of motion are at second-order, we evaluate the three-point correlation function of primordial scalar perturbations generated during inflation. We show that the shape of non-Gaussianities is well approximated by the equilateral type. The equilateral non-linear parameter $f_{\\rm NL}^{\\rm equil}$ is derived on the quasi de Sitter background where the slow-variation parameters are much smaller than unity. We apply our formula for $f_{\\rm NL}^{\\rm equil}$ to a number of single-field models of inflation--such as k-inflation, k-inflation with Galileon terms, potential-driven Galileon inflation, nonminimal coupling models (including field-derivative coupling models), and Gauss-Bonnet gravity. ",
        "introduction": "The inflationary paradigm \\cite{inflation} can successfully account for the observed temperature anisotropies in Cosmic Microwave Background (CMB) \\cite{WMAP7} as well as the galaxy power spectrum \\cite{LSS}. Although we have not identified the origin of inflation yet, the upcoming observations such as PLANCK \\cite{PLANCK} are expected to provide more high-precision data for the primordial scalar/tensor power spectra. This will allow us to distinguish between a host of theoretical models of inflation \\cite{inflationmodels} observationally.  In addition to the spectral index of scalar density perturbations $n_{\\cal R}$ and the tensor-to-scalar ratio $r$, the information of primordial scalar non-Gaussianities is useful to break the degeneracy between different models \\cite{Salopek,Gangui,Verde,KSpergel}. In standard slow-roll inflation in which cosmic acceleration is driven by the potential energy of a scalar field $\\phi$, the non-linear parameter $f_{\\rm NL}^{\\rm equil}$ characterizing the equilateral shape of non-Gaussianities is small ($|f_{\\rm NL}^{\\rm equil}| \\ll 1$) \\cite{Bartolo,Maldacena} (see also Refs.~\\cite{Rigo,Lyth,Cremi}). Meanwhile, in the presence of a non-linear field kinetic energy in $X=-\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi/2$ such as k-inflation \\cite{kinflation}, Dirac-Born-Infeld (DBI) models \\cite{DBI,DBI2}, (dilatonic) ghost condensate \\cite{ghost,ghost2}, and Galileon inflation \\cite{Galileoninf,Galileoninf2}, it is possible to realize large non-Gaussianities detectable in future observations ($|f_{\\rm NL}^{\\rm equil}| \\gtrsim 10$) \\cite{Seery,Gruzinov,Chen,Ohashi}.  The viable theoretical models of inflation are usually constructed to keep the equations of motion at second-order in order to avoid the Ostrogradski's instability \\cite{Ost}. In 1974 Horndeski \\cite{Horndeski} derived the most general single-field Lagrangian giving rise to second-order field equations. Recently there has been renewed interest in second-order gravitational theories in connection with the Dvali-Gabadadze-Porrati braneworld \\cite{DGP} and Galileon gravity \\cite{Nicolis,CGalileon} (see also Refs.~\\cite{Rham,genesis,Kamada,Naruko,Van,Trodden,petel,Elliston,Gao,Galileondark}). Deffayet {\\it et al.} \\cite{DGSZ} showed that the most general action in those theories is given by \\begin{equation} S = \\int d^{4}x\\sqrt{-g} \\biggl[\\frac{\\Mpl^{2}}{2}\\, R +P(\\phi,X)-G_{3}(\\phi,X)\\,\\Box\\phi +\\mathcal{L}_{4}+\\mathcal{L}_{5} \\biggr]. \\label{action} \\end{equation} Here $g$ is the determinant of the metric $g_{\\mu \\nu}$, $M_{\\rm pl}$ is the reduced Planck mass, $R$ is a Ricci scalar, and \\begin{eqnarray} \\mathcal{L}_{4} & = & G_{4}(\\phi,X)\\, R+G_{4,X}\\,[(\\Box\\phi)^{2} -(\\nabla_{\\mu}\\nabla_{\\nu}\\phi)\\,(\\nabla^{\\mu}\\nabla^{\\nu}\\phi)]\\,,\\\\ \\mathcal{L}_{5} & = & G_{5}(\\phi,X)\\, G_{\\mu\\nu}\\,(\\nabla^{\\mu}\\nabla^{\\nu}\\phi) -\\frac{1}{6}G_{5,X}[(\\Box\\phi)^{3} -3(\\Box\\phi)\\,(\\nabla_{\\mu}\\nabla_{\\nu}\\phi)\\,(\\nabla^{\\mu}\\nabla^{\\nu}\\phi) +2(\\nabla^{\\mu}\\nabla_{\\alpha}\\phi)\\,(\\nabla^{\\alpha}\\nabla_{\\beta}\\phi)\\, (\\nabla^{\\beta}\\nabla_{\\mu}\\phi)], \\end{eqnarray} where $G_i$ ($i=3, 4, 5$) are functions in terms of $\\phi$ and $X=-\\partial^{\\mu} \\phi \\partial_{\\mu} \\phi/2$ with the partial derivatives $G_{i,X} \\equiv \\partial G_i/\\partial X$, and $G_{\\mu \\nu}=R_{\\mu \\nu}-g_{\\mu \\nu}R/2$ is the Einstein tensor ($R_{\\mu \\nu}$ is the Ricci tensor). The action (\\ref{action}) is equivalent to that derived by Hordenski \\cite{Char,KYY}. In Ref.~\\cite{KYY} the spectra of scalar and tensor perturbations were derived for the above models.  In this paper we compute the three-point correlation function of scalar metric perturbations for the inflationary models described by the action (\\ref{action}). Note that the calculation of scalar non-Gaussianities up to the term $G_3 (\\phi,X) \\square \\phi$ was carried out in Refs.~\\cite{nongaugali,nongaugali2,nongaugali3} (see also Ref.~\\cite{Qiu} for non-minimal coupling models with $F(\\phi)R$). Taking into account the Lagrangians ${\\cal L}_4$ and ${\\cal L}_5$, we can cover a wide variety of gravitational theories such as scalar-tensor theories \\cite{BD} ($G_4=F(\\phi)$), field derivative couplings with gravity \\cite{Germani,Germani2} ($G_5=G(\\phi)$), Galileon gravity \\cite{Nicolis,CGalileon} ($G_3 \\propto X$, $G_4 \\propto X^2$, $G_5 \\propto X^2$), higher-curvature gravity \\cite{fRreview} (including Gauss-Bonnet and $f(R)$ theories), $\\alpha'$ corrections appearing in low-energy effective string theory \\cite{Tseytlin}.  Even in the presence of the new contributions coming from the terms $G_4$ and $G_5$ we shall show that the shape of non-Gaussianities can be well approximated by the equilateral type. The equilateral non-linear parameter $f_{\\rm NL}^{\\rm equil}$ is derived on the quasi de Sitter background. We also apply our formula to concrete inflationary models in order to present the cases in which large non-Gaussianities can be realized.  This paper is organized as follows. In Sec.~\\ref{modelsec} we derive the background equations of motion on the flat Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) background and introduce a number of ``slow-variation'' parameters convenient for the calculation of primordial non-Gaussianities. In Sec.~\\ref{sactionsec} the action (\\ref{action}) is expanded at second-order in perturbations in order to obtain the power spectrum of curvature perturbations generated during inflation. We also obtain the spectrum of gravitational waves and the tensor-to-scalar ratio. In Sec.~\\ref{thirdsec} we expand the action (\\ref{action}) at third-order in perturbations and compute the three-point correlation function of curvature perturbations. In Sec.~\\ref{shapesec} we study the shape of non-Gaussianities for the new terms appearing in the three-point correlation function. In Sec.~\\ref{quasisec} we derive an explicit formula for the equilateral non-linear parameter $f_{\\rm NL}^{\\rm equil}$ under the slow-variation approximation. In Sec.~\\ref{applysec} our formula for $f_{\\rm NL}^{\\rm equil}$ is applied to a number of concrete models of inflation. Sec.~\\ref{concludesec} is devoted to conclusions. ",
        "conclusions": "\\label{concludesec}  In this paper we have evaluated the primordial non-Gaussianities generated during inflation for the models described by the action (\\ref{action}). Our analysis is general enough in that it covers a wide variety of single-field models having second-order equations of motion.  The procedure to obtain the three-point correlation function of curvature perturbations is similar to that given in Ref.~\\cite{nongaugali2}. There are 8 shape functions which contribute to the scalar non-Gaussianities. Among them, five functions are already known to give rise to the equilateral type of non-Gaussianities. We studied the shapes of the remaining three functions and found that they can be well approximated by the same type as well. Hence the dominant contribution to the three-point correlation function comes from the equilateral configuration.  We derived the equilateral non-linear parameter $f_{\\rm NL}^{\\rm equil}$ under the approximation that the slow-variation parameters defined in Eq.~(\\ref{slowva}) are much smaller than 1. The formula (\\ref{fnlformula}) is valid for any quasi de Sitter background and thus it is convenient to apply it to concrete single-field models of inflation. In fact we used this formula to a number of models such as (A) k-inflation, (B) k-inflation with the terms $G_i(X)$ ($i=3,4,5$), (C) potential-driven Galileon inflation, (D) nonminimal coupling models without a non-linear term in $X$, (E) potential-driven inflation with a Gauss-Bonnet term. In the models (D) and (E) we showed that $c_s^2$ is close to 1 and hence the non-Gaussianities are small. However, for the models (A), (B), (C), it is possible to realize $|f_{\\rm NL}^{\\rm equil}| \\gg 1$ depending on the choice of the functions $P, G_i$ ($i=3,4,5$).  The potential detectability of non-Gaussianities in future observations will allow us to distinguish between different inflationary models. In particular we expect that the precise measurement of $f_{\\rm NL}^{\\rm equil}$ as well as $n_{\\cal R}$ and $r$ will provide significant implications for the viability of single-field models in which $c_s^2$ is much smaller than 1."
    },
    "1107/1107.1253_arXiv.txt": {
        "abstract": "Observations of the most distant bright quasars imply that billion solar mass supermassive black holes (SMBH) have to be assembled within the first eight hundred million years. Under our standard galaxy formation scenario such fast growth implies large gas densities providing sustained accretion at critical or supercritical rates onto an initial black hole seed. It has been a long standing question whether and how such high black hole accretion rates can be achieved and sustained at the centers of early galaxies.  Here we use our new {\\it MassiveBlack} cosmological hydrodynamic simulation covering a volume $(0.75 \\Gpc)^3$ appropriate for studying the rare first quasars to show that steady high density cold gas flows responsible for assembling the first galaxies produce the high gas densities that lead to sustained critical accretion rates and hence rapid growth commensurate with the existence of $\\sim 10^9 \\Msun$ black holes as early as $z\\sim 7$. We find that under these conditions quasar feedback is not effective at stopping the cold gas from penetrating the central regions and hence cannot quench the accretion until the host galaxy reaches $M_{\\rm halo} \\simgt 10^{12} \\Msun$. This cold-flow driven scenario for the formation of quasars implies that they should be ubiquitous in galaxies in the early universe and that major (proto)galaxy mergers are not a requirement for efficient fuel supply and growth, particularly for the earliest SMBHs. ",
        "introduction": "} It is now well established that the properties of supermassive black holes (SMBH) found at the centers of galaxies today are tightly coupled to those of their hosts implying a strong link between black hole and galaxy formation. The strongest direct constraint on the high-redshift evolution of SMBHs comes from the observations of the luminous quasars at $z \\sim 6$ in the Sloan Digital Sky Survey (SDSS) \\citep{Fan2006, Jiang2009} and even more recently at $z=7$ \\citep{Mortlock2011}.  Although rare (the comoving space density of $z \\sim 6$ quasars is roughly $n \\sim$ a few$ \\Gpc^{-3}$) the inferred hole masses of these quasars are in excess of $10^9~\\Msun$ comparable to the masses of the most massive black holes in the Universe today. The origin of these massive black hole seed and the physical conditions that allow early growth to supermassive black holes remain a challenging problem.  \\begin{figure*} \\centering \\hbox{ \\hspace{-0.1cm} \\includegraphics[scale=1.02]{figs/tdm-paper.ps}} \\label{fullsim} \\caption{The cosmological mass distribution in our the simulation volume at $z=5$. The projected gas density over the whole volume ('unwrapped' into 2D) is shown in the large scale image. The two overlaid panels show successive zoom-ins by factor of 10, center on the region where the most massive black hole is found.} \\end{figure*} In order to have had sufficient time to build up via gas accretion and BH mergers (resulting from the hierarchical merging of their host halos) the first 'seed' black holes must have appeared at early epoch, $z > 10$.  The origin and nature of this seed population remain uncertain. Two distinct populations of seed masses, in the range of $100-10^5~\\msun$ have been proposed: the small mass seeds are usually thought to be the remnants of the first generation of PopIII stars formed of metal-free gas at $z\\sim 20-30$ \\citep[e.g.]{Bromm1999, Abel2000, Nakamura2001, Yoshida2003, Gao2006}, while the large seeds form in direct dynamical collapse in metal-free galaxies \\citep[although see also Mayer et al. 2010 for direct collapse into a massive blackhole in metal enriched regime]{Koushiappas2004, Begelman2006}.  Growing the seeds to $~10^9 \\Msun$ in less than a billion years requires extremely large accretion rates - as mergers between black holes are too rare and too inefficient for significant growth.  For a black hole accreting at the critical Eddington accretion rate its luminosity $L_{\\rm Edd} = (4 \\pi G c m_p)/ \\sigma_{T} M_{BH} = \\eta \\Mdot_{\\rm Edd} c^2$ (where $G$, $c$, $m_{p}$ and $\\sigma_{T}$ are the gravitational constant, speed of light, proton mass and Thomson cross section, and $\\eta$ is the standard accretion efficiency) implies an exponential growth at the characteristic Eddington timescales, $t_{\\rm Edd} = 450 \\eta / (1- \\eta)$ Myr, such that $M_{\\rm BH} = M_{\\rm seed} e^{t/t_{\\rm Edd}}$. For a seed mass ranging from $M_{\\rm seed} \\sim 100-10^5 \\Msun$ this requires $10-17$ e-foldings to reach $M_{\\rm BH} \\sim 10^9 \\Msun$. The crucial question (for any given seed model) is then where (if at all, in which kind of halos) and how (at what gas inflow rates) such vigorous accretion can be sustained at these early times.  As bright quasars are likely to occur in extremely rare high-density peaks in the early universe, large computational volumes are needed to study them. Here we use a new large cosmological Smooth Particle hydrodynamics (SPH) simulation, {\\it Massive Black} (covering a volume of $[0.75 \\Gpc]^{3}$) with sufficiently high resolution (over 65 billion particles) to be able to include tested prescriptions for star formation, black hole accretion and associated feedback processes to investigate whether and if such objects may be formed within our standard structure formation models.  Crucially, our {\\it MassiveBlack} simulation is of sufficiently high-resolution to allow us to follow the mass distribution in the inner regions of galaxies and hence model star formation and black hole growth directly and self-consistently while still evolving a close to Gigaparsec scale region. It therefore provides a unique framework to study the formation of the first quasars. ",
        "conclusions": "With our new large cosmological simulation {\\it MassiveBlack} we show that the short timescale associated with infall via cold flows and the short cooling timescales in cold radial streams that penetrate the halo render the flow into the central regions unstoppable by feedback allowing it to easily sustain BH growth at the Eddington rates to build up the required BH masses by $z=6-7$. One consequence of this scenario is that BH masses at these redshifts are expected to show deviations from the local $M_{BH} - \\sigma$ relation. BH masses assembled faster (most growth occurs over few hundred million years) than the stellar spheroid (assembled over Gyrs timescales).  Black hole masses larger than that inferred from the local $M_{BH} - \\sigma$ relation have infact been suggested for the first quasars.  Stream fed accretion will still be relevant over the peak of the quasar phase but mergers (as merger rates peak closer to those redshift) will become an increasingly major player in their growth and formation. We speculate however that most of the growth of a quasar's mass is likely to always occur before the shock heating scale of an halo is reached much like most of its star formation rate \\citep{Dekel2009}. We will investigate this further in our large volume in future work.  Our scenario is somewhat similar to that proposed by \\citet{Mayer2010} or \\citet{Li2007} yet, crucially it does not rely on a major merger to induce the strong gas inflows but points to a more common origin for them particularly at these redshifts (which we could find by virtue of having a large volume in our simulation, see also \\citet{Sijacki2009} who followed the build-up of a single hihigh-z quasar). As we have shown, relaxing the constraint for massive mergers makes it plausible to attain quite commonly large black hole masses as high as $z=7$ commensurate with \\citet{Mortlock2011}."
    },
    "1107/1107.4640_arXiv.txt": {
        "abstract": "We present Gran-Telescopio-Canarias/OSIRIS optical spectra of 4 of the most compact and massive early-type galaxies in the Groth Strip Survey at redshift $z\\sim1$, with effective radii $\\Reff=0.5 - 2.4$ kpc and photometric stellar masses $\\Mstar=1.2 - 4\\times10^{11} \\Msun$. We find these galaxies have velocity dispersions $\\sigma = 156 - 236$ \\kms. The spectra are well fitted by single stellar population models with approximately 1 Gyr of age and solar metallicity. We find that: i) the dynamical masses of these galaxies are systematically smaller by a factor of $\\sim$6 than the published stellar masses using $BRIJK$ photometry; ii) when estimating stellar masses as 0.7$\\times$\\Mdyn, a combination of passive luminosity fading with mass/size growth due to minor mergers can plausibly evolve our objects to match the properties of the local population of early-type galaxies. ",
        "introduction": "The evolution of the structural properties of massive early-type galaxies (ETGs) is currently a topic of debate with many open questions. In recent years, studies have shown that galaxies with $\\Mstar \\geq 10^{11}\\Msun$ are smaller by a factor of four at $\\z \\geq1.5$ \\citep{daddi05,trujillo06,longhetti07,buitrago08,damjanov09} and a factor of two at $\\z \\sim1$ \\citep{trujillo07, trujillo11} than the nearby population of the same mass. For the most compact and massive galaxies in particular, \\citet{trujillo09} found that only $\\sim$0.03\\% of SDSS galaxies with $\\Mstar \\geq8\\times10^{10} \\Msun$ at $\\z \\leq0.2$ are ETGs with $\\Reff \\leq1.5$ kpc. This contrasts with the much higher fraction $\\sim$15\\% of such galaxies found at \\z$\\sim$1 by Trujillo et al. (2007, hereafter T07). Most interestingly, the luminosity-weighted ages of the local compact ETGs from \\citet{trujillo09} are very low ($\\sim$2 Gyr), which argues against this sample being the counterpart of a passively evolved population of high redshift galaxies. Thus, this suggests that most high redshift compact and massive galaxies must have undergone a systematic growth in size.\\\\ Among the different mechanisms proposed to explain this expansion, the one that reproduces best the observed evolution of the average mass-size relation \\citep{trujillo11} is the minor merger scenario \\citep{naab09,hopkins10}. With this mechanism, galaxies grow inside-out by accretion of smaller satellites that build up an extended envelope \\citep{dokkum10,hopkins09}. This model predicts a substantial growth in mass and effective radius, with a mild decrease in stellar velocity dispersion $\\sigma$. Thus, measurements of $\\sigma$ at different redshifts are a direct way to constrain this evolutionary scenario. In addition, they allow for an independent mass estimate through the virial theorem.\\\\ The minor merger scenario is consistent with the common findings of $\\sigma \\lesssim 250$ \\kms for general ETG samples at $\\z=1-2$ \\citep{gebhardt03,treu05,wel05,serego05,cenarro09,cappellari09,fernandezlorenzo11}. Conversely, it is greatly challenged by the dramatic evolution required for the super-dense galaxies ($\\Reff<1$ kpc, $\\Mstar \\geq10^{11}\\Msun$), which are expected to have much larger $\\sigma$. \\citet{dokkum09} measured  $\\sigma \\simeq 500\\pm100$ \\kms from a super-dense ETG at $z=2.2$, yielding a dynamical mass that is in agreement with the photometric mass estimate. However, more $\\sigma$ measurements of these extreme galaxies are needed to support or weaken the need of a stronger evolutionary scenario. In this paper we present $\\sigma$'s of four massive and compact ETGs at $\\z \\sim1$, two of which are super-dense. We compare dynamical/stellar masses and estimate the evolutionary paths of these objects. Throughout this work we adopt a standard cosmology of $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$ \\kms Mpc$^{-1}$.    \\begin{figure}[h] \\includegraphics[scale=.62]{fig1.eps} \\caption{\\label{fig1}Spectra of our sample galaxies. The black solid lines represent the MOVEL fitting results. Light grey regions are [OII] emission at $3727\\text{ \\AA}$ and telluric absorption at $7600\\text{ \\AA}$, which were excluded from the fits (see text in \\S 2.2). } \\end{figure}  \\begin{figure}[h] \\includegraphics[scale=.77]{fig2.eps} \\caption{\\label{fig2}\\textbf{a)} Stellar mass-size relation at $z\\sim1$. Small circles represent the entire sample from T07, black squares are our selection among them. The straight line is the local relation from \\citet{shen03}. Black diamonds are our galaxies with stellar masses derived as $0.7\\,M_{\\text{dyn}}$. Gray diamonds are their position when evolved to $z=0$ via minor merger growth. \\textbf{b)} B-band Fundamental Plane of ETGs. The solid line is the local relation from \\citet{jorgensen96}. Our sample is shown as black diamonds. Gray diamonds are their position when evolved to $z=0$ via minor merger growth and passive luminosity fading. Samples from \\citet{gebhardt03} and \\citet{wel05} in the range $z=0.8-1.1$ are shown for comparison. We transformed our $I$ to $B$ magnitudes appling \\ensuremath{k}-corrections based on $V-I$ color as in \\citet{gebhardt03}. In addition, we applied a negative offset of 0.1 mag on all surface brightness data from \\citet{gebhardt03} to transform their Vega to our AB magnitude system.  } \\end{figure} ",
        "conclusions": "Dynamical masses are calculated under the assumption of homology as $M_{\\mathrm{dyn}}=\\beta \\Reff \\sigma^2/G$ with $\\beta=5$ (see Table 1), following studies of local ETGs \\citep{cappellari06}. As noted in $\\S$ 2.2, the galaxy 12024790 shows a double peak in the [OII] emission. If such feature is due to rotation of a gaseous disk, it yields a projected velocity of $V_r \\mathrm{sin}(i)\\simeq165$ \\kms, which is comparable to its measured $\\sigma=156$ \\kms. However, it is unlikely that this caused an important bias in the dynamical mass estimates. At $\\z=0$ it has been found that $\\beta=5$ is robust against rotation up to $V_r/\\sigma \\approx 1$ \\citep{cappellari06, cappellari07}, and there is no clear evidence that this would be different at $\\z=1$ \\citep{wel08b}. Therefore, we assume that any deviations from homology due to possible rotational support are within the errors of our dynamical masses. \\\\Our galaxies have values that are systematically smaller than the originally published stellar masses by an average factor of $\\sim$6. This is much larger than the combined error of photometric stellar mass and radius (typically a factor of $\\sim$3). Note that for the galaxy 12024321 these masses are consistent within uncertainties. The stellar population models we used with the RAINBOW data yield stellar masses in the range of log$\\Mstar=10.65-10.98$, on average a factor of $\\sim$1.8 larger than the dynamical masses. The masses from RAINBOW are derived with the same IMF and stellar population models as in T07 but with a broader spectral range (UV-FIR) as described in $\\S$ 2.4. In order to investigate the large differences between both sets of masses, we also derived masses using only $BRIJK$ bands as in T07. We found no significant difference between our UV-FIR and $BRIJK$ masses, implying that the discrepancy with T07 cannot be simply explained by the use of a different spectral range. We also checked that our mass determinations are weakly sensitive to the choice of extinction model and small variations in metallicity. \\citet{barro11b} derive masses for hundreds of galaxies in T07's sample using the full RAINBOW photometry. They find that, on average, their masses are consistent with those from T07. Therefore, the cause of the inconsistency with our photometric masses is unclear. Given the large uncertainties in the systematics of photometric masses, we adopt dynamical masses as the correct total mass estimates. Following local lensing studies \\citep{gavazzi07}, we apply a fiducial dark matter fraction $f_{\\mathrm{DM}}=0.3$ and recalculate the stellar masses as $0.7\\times M_{\\mathrm{dyn}}$. \\\\Comparison studies between stellar and dynamical mass of less dense ETGs than those in our sample have been conducted at low \\citep{cappellari06} and high redshifts \\citep{wel06} with the result of dynamical masses being on average equal or larger than stellar masses, as expected when considering a dark matter contribution. When \\citet{wel06} calculate stellar masses using similar multi-band photometry and stellar population parameters as in T07, they find $M_\\star/M_{\\mathrm{dyn}}=0.66$ (Salpeter IMF). Therefore, our results point towards a systematic overestimation of the stellar masses used to select the galaxies of our sample. We cannot rule out a deviation from homology that requires $\\beta \\neq 5$. However, the theoretical expectation for objects with such high S\\'ersic index is $\\beta \\lesssim 5$ \\citep{cappellari06,bertin02}, which would yield even smaller dynamical masses. \\\\Fig. \\ref{fig2}a shows the position of our objects in the stellar mass-size relation with the old/new stellar mass estimates. Fig. \\ref{fig2}b shows our objects in the Fundamental Plane (FP). At a given $\\Reff$, they have systematically larger masses and smaller combination of surface brightness (SB$_e$) and $\\sigma$ than the respective local relations. Moreover, our objects follow a trend in which they lie further from both local relations at smaller radii. This result has also been found in $\\z\\sim0.9$ clusters by \\citet{jorgensen06,jorgensen07}. Compared to the FP derived by these authors, at a given $\\Reff$ our field sample shows a similar slope and a relative offset of $\\sim0.3$ mag towards brighter $\\mathrm{SB}_e$, which is in agreement with field galaxies having generally younger stellar populations than clusters of the same redshift. \\\\In order to understand how our galaxies relate to low redshift ETGs, we investigate the simple evolutionary scenario based on passive luminosity fading of the stellar populations and growth via minor mergers that has been proposed for the general population of ETGs at $\\z \\sim 1$ \\citep{gebhardt03, naab09}. To derive the amount of fading for our sample, we use \\citet{bc03} models with Chabrier IMF and the redshift, ages and metallicities from Table \\ref{table1}. Passive evolution from $\\z \\sim1$ to $\\z=0$ results in an average dimming of $\\sim$2 mag in SB$_e$. We assume that the passive luminosity evolution does not affect $\\Reff$. Minor mergers are very efficient growing the size at a given mass increment, slightly decreasing $\\sigma$. According to the simulations by \\citet{naab09}, an ETG with $M_\\star =10^{11}M_\\odot$ and $\\Reff=1.5$ kpc would grow a factor of 1.5 in mass via minor mergers from $\\z \\sim1$ to $\\z=0$. We adopt this factor of mass growth for our galaxies. Following these authors, we model each merger as $\\sigma^2_f/\\sigma^2_i=(1+\\eta \\epsilon)/(1+\\eta)$ and ${\\Reff}_f/{\\Reff}_i=(1+\\eta)^{2}/(1+\\eta \\epsilon)$, where the subindexes (i)/(f) denote initial/final values of the host galaxy and $\\epsilon=\\sigma^2_a/\\sigma^2_i$ and the mass ratio $\\eta=M_a/M_i$ take into account the accreted system ($M_a,\\sigma_a$). We assume a history of four mergers with constant $\\eta=0.1$ and $\\epsilon=0.2$ for all galaxies (note that $\\eta$ and the total mass growth set the total number of mergers). This choice of values is consistent with the simulations by \\citet{naab09}. The total change for $\\Reff$ and $\\sigma$ become factors of 2 and 0.8, respectively. These results are robust against variations in the values of $\\eta$ and $\\epsilon$ at the given total mass growth. We assume that the mass-to-light ratio does not change significantly after each merger. In addition, we assume that as the size increases, SB$_e$ decreases by $\\Delta \\mathrm{SB}_e\\propto2.94\\Delta \\mathrm{log}\\Reff$ \\citep{hamabe87}. \\\\Fig. \\ref{fig2} shows the combined evolution of minor mergers and luminosity fading for our galaxies. Most of them reach positions that are consistent with the local FP and mass-size relations, considering the estimated uncertainties in our measurements. Therefore, this simple evolutionary scenario can plausibly describe the evolution in our sample galaxies to become normal ETGs at $\\z=0$."
    },
    "1107/1107.3898_arXiv.txt": {
        "abstract": "Using an optically-unbiased selection process based on the HIPASS neutral hydrogen survey, we have selected a sample of 83 spatially isolated, gas-rich dwarf galaxies in the southern hemisphere with $cz$ between 350 and 1650 km s$^{-1}$, and with R-band luminosities and H~\\textsc{i} masses less than that of the Small Magellanic Cloud. The sample is an important population of dwarf galaxies in the local Universe, all with ongoing star formation, and most of which have no existing spectroscopic data. We are measuring the chemical abundances of these galaxies, using the Integral Field Spectrograph on the ANU 2.3m telescope, the Wide Field Spectrograph (WiFeS). This paper describes our survey criteria and procedures, lists the survey sample, and reports on initial observations. ",
        "introduction": "The main epoch of galaxy formation, occurring at redshifts $1 \\lesssim z \\lesssim 4$ (e.g. \\citeauthor{mad98} \\citeyear{mad98}) is long past. Large galaxies have evolved far from the pristine state in which they formed. Mergers, galactic winds, cycles of star birth and death and outflows from active galactic nuclei have not only chemically enriched the interstellar medium of these galaxies, but have also chemically enriched their surrounding inter-galactic space \\citep{kob07}. Theory predicts that chemical abundances will rise as newly-collapsing galaxies evolve and undergo successive generations of star formation. However, the spread of abundances should narrow with time as the interstellar gas becomes either internally mixed, diluted with infalling chemically pristine gas, or depleted by starburst-driven galactic-scale outflows \\citep{del04,kob07}. Thus, the chemical abundance in galaxies is inextricably linked to star formation and galaxy growth, providing a record of the generations of cosmic star formation, mass accretion, and mass-loss in galaxies.  In this regard, the gas-rich dwarf galaxies are of particular interest. These galaxies have processed a much smaller fraction of their gas through stars, and so are much less chemically evolved than the more massive disk or elliptical galaxies. Many of them have formed and developed in ``quiescent'' regions of space, possibly at very early times, away from the centres of dense clusters and the gravitational harassment of  massive neighbors, so their star formation and chemical evolution history may well be simpler than for the more massive systems.  A detailed study of the stellar and gas content of these gas-rich dwarf galaxies may therefore be expected to provide answers to the following key questions: \\begin{itemize} \\item{What is the relationship between mass and chemical abundance for low mass galaxies? The answer could provide sensitive constraints on the mass fraction of the interstellar medium lost to inter-galactic space by galaxy winds. We would like to know whether matter lost from dwarf galaxies can account for the ``missing baryon'' problem---the discrepancy between the baryonic mass currently contained in visible galaxies, and the baryonic mass inferred from the standard model of cosmology \\citep{bre07}.}  \\item{What is the total mass of oxygen in each dwarf galaxy? This is determined from the chemical abundance and the total H~\\textsc{i} mass of these galaxies. Since we know how many stars have formed in the galaxy from I-band photometry, we may obtain tight constraints on the fraction of heavy elements that have been lost to the inter-galactic medium by galactic winds in these systems.}  \\item{Is there a chemical abundance floor in the local Universe? Such a floor, predicted to be about 1/100 solar \\citep{kob07}, would result from pollution of inter-galactic gas by heavy elements ejected in starburst-powered galactic winds, or by black hole jet-powered outflows from massive galaxies.} \\end{itemize}  There is evidence that dwarf glaxies have a wide variety of evolutionary histories (e.g. \\citet{gre97,mat98,tol09}) and that---possibly as a result of this---they display a wide scatter on the mass-metallicity relationship \\citep{tre04, lee06, gus09}. However, the observed sample of objects in this region of the parameter space is still relatively small.  Although these dwarf systems are dominated by dark matter \\citep{mat98} and they retain much of their original gas content, the inferred chemical yields are much lower than those estimated for more massive galaxies. This suggests that these galaxies may have formed very early in the history of the Universe and have subsequently had quite severe episodes of mass-loss through galactic winds. The study of small isolated gas-rich dwarf galaxies may provide evidence of the conditions and processes long since erased in larger galaxies and clusters.  Dwarf irregular galaxies (`dwarf' defined for the purposes of this study as having a gas+stellar mass less than that of the Small Magellanic Cloud) are numerous throughout the Local Volume but they remain surprisingly poorly studied. Optical catalogs by their nature tend to be incomplete and affected by the Malmquist Bias \\citep{mal21}.  Dwarf galaxies can be missed due to their low surface brightness, or mistaken for much more massive and distant objects. The SINGG survey  \\citep{meu06} went a long way towards rectifying this situation as it drew its objects from the HIPASS neutral hydrogen survey \\citep{zwa04, mey04, kor04}. This provided a complete volume-limited sample for a given H~\\textsc{i} mass content. Follow-up R-band and H$\\alpha$ photometry revealed both the stellar luminosity and the star-formation rate in these galaxies. Much to the surprise of \\citet{meu06}, evidence for on-going ($\\lesssim 10$~Myr) star formation was found in nearly every case. This opens up the possibility of follow-up spectroscopy to establish the chemical abundances in the H~\\textsc{ii} regions of these objects.  Going beyond the SINGG survey, which investigated only $\\sim$ 10\\% of the HIPASS sources, there are many other small isolated gas-rich dwarf galaxies in the HOPCAT catalogue \\citep{doy05}, which could also provide insight into dwarf galaxy evolution. Both the SINGG and HOPCAT galaxies can now be studied very efficiently using integral field units (IFUs). In particular the Wide-Field Spectrograph (WiFeS; Dopita et al. 2007, 2010) at the Australian National University (ANU) 2.3m telescope at Siding Spring Observatory is ideally suited for the study of these galaxies. The WiFeS instrument is a highly efficient double-beam, image-slicing integral-field spectrograph with spectral resolutions $R=3000$ and 7000, covering the wavelength range 3500\\AA~ to $\\sim$9000\\AA.  It offers a contiguous $25 \\times 38$~arc sec field-of-view, well matched to the angular size of these objects. Such an instrument obviates the need to obtain optical photometry in either broad or narrow bands, and allows us to extract the complete spectra of each of the individual  H~\\textsc{ii}  regions present in the galaxy.  Motivated by the opportunity to study dwarf galaxy formation and evolution in the Local Volume (D $\\lesssim$ 20 Mpc), we have identified a volume-limited sample of small, isolated gas-rich dwarfs. In this paper, we describe the characteristics of this sample, and describe initial results. ",
        "conclusions": "We present a sample of 83 small isolated gas-rich dwarf irregular galaxies identified using their neutral hydrogen 21cm signatures and presnece of star formation.  The sample consists of galaxies with lower neutral hydrogen masses and lower R-band luminosities than the SMC. They are located in the southern sky with heliocentric recession velocities between 350 and 1650 km s$^{-1}$.  We are observing these objects using the WiFeS integral field spectrograph on the ANU 2.3m telescope at Siding Spring, to measure nebular metallicities.  We intend to use this data  to explore the mass-metallicity relation at the low mass end of the spectrum, and to see if there is evidence for a metallicty floor in the IGM. The sample appears to include a higher percentage of Blue Compact Dwarf Galaxies than expected from other optical surveys, and a number of unusual extended dwarf objects. \\\\ \\\\"
    },
    "1107/1107.2251.txt": {
        "abstract": "This study, using publicly available simulations, focuses on the characterisation of the non-Gaussianity produced by radio point sources and by infrared (IR) sources in the frequency range of the cosmic microwave background from 30 to 350 GHz.  We propose a simple prescription to infer the angular bispectrum from the power spectrum of point sources considering independent populations of sources, with or without clustering. We test the accuracy of our prediction using publicly available all-sky simulations of radio and IR sources and find very good agreement.  We further characterise the configuration dependence and the frequency behaviour of the IR and radio bispectra. We show that the IR angular bispectrum peaks for squeezed triangles and that the clustering of IR sources enhances the bispectrum values by several orders of magnitude at scales $\\ell \\sim 100$.  At 150 GHz the bispectrum of IR sources starts to dominate that of radio sources on large angular scales, and it dominates over the whole multipole range at 350 GHz.  Finally, we compute the bias on $f_\\mathrm{NL}$ induced by radio and IR sources. We show that the positive bias induced by radio sources is significantly reduced by masking the sources.  We also show, for the first time, that the form of the IR bispectrum mimics a primordial `local' bispectrum $f_\\mathrm{NL}$. The IR sources produce a negative bias which becomes important for Planck-like resolution and at high frequencies ($\\Delta f_\\mathrm{NL} \\sim -6$ at 277 GHz and $\\Delta f_\\mathrm{NL} \\sim -$60-70 at 350 GHz). Most of the signal being due to the clustering of faint IR sources, the bias $\\Delta f_\\mathrm{NL}^\\mathrm{IR}$ is not reduced by masking sources above a flux limit and may, in some cases, even be increased due to the reduction of the shot-noise term. ",
        "introduction": "In the last few decades, the Cosmic Microwave Background (CMB) has become a very successful probe of the early and late universe. The smallness of the perturbations in the cosmic fluids, and hence in the space-time metric, allows us to use linear perturbation theory to compute their evolution efficiently and accurately through Boltzmann codes (\\cite{Seljak1996}, \\cite{Lewis2000}, \\cite{Lesgourgues2011}).  Since COBE \\citep{Smoot1992} the measurement of the CMB power spectrum has been achieved by many experiments and over a wide range of scales. The most recent CMB data are those of the Seven-Year Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{Larson2011}, Atacama Cosmology Telescope (ACT) \\citep{Das2011} and South Pole Telescope (SPT) \\citep{Keisler2011}. Constraints from all these measurements, combined with probes of the geometry of the Universe , like Baryonic Acoustic Oscillations (e.g. \\cite{Percival2010}, \\cite{Blake2011}), type Ia Supernovae (e.g. \\cite{Astier2006}, \\cite{Hicken2009}, \\cite{Guy2010}) and Hubble constant measurements (e.g. \\cite{Freedman2001}, \\cite{Riess2009}, \\cite{Freedman2009}), give a converging view of our Universe and have lead to the establishment of a `standard model' of cosmology (e.g. \\cite{Larson2011}) known as $\\Lambda$CDM. In this model, the universe is flat, dominated by a cold dark matter component and a `dark energy' component compatible with a cosmological constant. The present constraints suggest that CMB anisotropies are a realisation of a primordial random process that generated the initial perturbations from quantum fluctuations which were then stretched to macroscopic scales by inflation (e.g. \\cite{Starobinski1979}, \\cite{Guth1981}, \\cite{Liddle2000}, see \\cite{Bassett2006} and \\cite{Linde2008} for reviews).  The microwave sky is however not made of CMB primary signal alone, it consists also of secondary anisotropies such as those associated with the Integrated Sachs-Wolfe effect -- arising from the time-variable metric perturbations \\citep{SW1967} --, those arising from the Sunyaev-Zel'dovich (SZ) effect (inverse Compton scattering) in the direction of galaxy clusters \\citep{Sunyaev1972}, those due to the Doppler effect from moving structures (e.g. kinetic SZ effect from clusters and inhomogeneous reionisation), see \\cite{Aghanim2008} for a review on secondary anisotropies. In addition, there are other astrophysical components in the microwave domain such as the dust, synchrotron and free-free emissions from our Galaxy \\citep{Planck-Collaboration-Dust}, the emission from radio sources that dominate at lower frequencies but contribute significantly at microwave frequencies \\citep{deZotti2005}, and the emission from dusty star-forming galaxies emitting mainly in the Infra-Red (IR) domain \\citep{Low1968}.  In this study, we will focus only on the characterisation of the extra-galactic point sources, namely the radio sources and the IR dusty galaxies. They contribute notably to the power spectrum at CMB frequencies and start dominating over the CMB at about $\\ell \\sim$ 2000 when the CMB signal plummets.  Active Galactic Nuclei (AGN) are observed as radio sources via their synchrotron emission. They have been widely studied in the CMB context especially at low frequencies $\\nu \\leq 90$ GHz (e.g. \\cite{deZotti2005}, \\cite{Boughn2008}, \\cite{Sajina2011}). They affect mostly the lower end of frequencies observed by CMB experiments. Their largest impact was thought to be in the frequency bands from 30 to 90 GHz but the recent Planck results \\citep{Planck-Collaboration-ERCSC, Planck-Collaboration-radioSED} show that radio sources are detected at frequencies as high as 217 GHz. At the CMB frequencies, radio sources do not cluster (\\cite{Toffolatti1998} \\cite{Gonzalez2005}) and thus exhibit a flat power-spectrum (see Appendix \\ref{appendix:shot}).  Star-forming galaxies are observed as IR sources via the thermal emission from dust heated by the ultra-violet emission of young stars. These IR sources contribute to the signal observed by CMB experiments at frequencies higher than 150 GHz, thus they are particularly relevant for the most recent CMB experiments, e.g. SPT, ACT and Planck High Frequency Instrument observations. The cumulated emission from the IR sources, the Cosmic Infrared Background (CIB), was first discovered by \\cite{Puget1996}, and its anisotropies were first characterised by \\cite{Lagache2000} and \\cite{Matsuhara2000}. Many other observations were possible in the last decade in the IR and submm domain (\\cite{Lagache2007}, \\cite{Viero2009}, \\cite{Hall2010}, \\cite{Amblard2011}, \\cite{Planck-Collaboration-CIB}). In particular, the latest constraints of the CIB from Planck showed that its power spectrum, at frequencies 217-353-545-857 GHz, is well modeled by a power law $C_\\ell^\\mathrm{CIB}(\\nu) = A(\\nu)\\times \\left(\\frac{\\ell}{1000}\\right)^n$ with e.g. $A=(104 \\pm 4)\\times 10^2 \\mathrm{Jy}^2/\\mathrm{sr}$ and $n=-1.08 \\pm 0.06$ at 545 GHz. This behaviour contrasts with the flat spectrum of radio sources and is due to the clustering of the IR galaxies and their host dark-matter halos.  These point sources are super-imposed on the primordial fluctuations. The simplest models of inflation (single-field and slow-roll) predict a small primordial non-Gaussianity (NG) (\\cite{Acquaviva2003}, \\cite{Maldacena2003}, \\cite{Creminelli2004})  that is sub-dominant to the NG induced by the non-linear evolution of the perturbations, a contribution that is necessarily always present. More complex inflationary models, e.g. multi-field scenarios, may predict larger NG \\citep{Byrnes2010}, to the point of being detectable. A simple and yet powerful probe of non-Gaussianity is the three-point function in harmonic space, the angular bispectrum (see section \\ref{sect:3pNGe} for more details), which is defined as a function of a multipole triplet ($\\ell_1,\\ell_2,\\ell_3$). The angular bispectrum vanishes for a Gaussian field like all odd-order moments. Besides the bispectrum, connected even-order moments may also be used to probe non-Gaussianity, and the 4-point function or trispectrum has indeed been a focus of interest \\citep{Kunz:2001ym}, especially for lensing studies \\citep{Cooray2002}.  There are many different models of inflation, and they often predict very similar power spectra that are close to being scale invariant. For models that lead to a measurable bispectrum, however, this degeneracy can often be broken by studying the dependence of the bispectrum amplitude on ($\\ell_1,\\ell_2,\\ell_3$), e.g. a large signal for squeezed triangles is indicative of slow-roll multi-field inflation models, equilateral triangles are enhanced for models with non-canonical kinetic terms, and folded triangles for non-standard vacuum initial conditions (e.g. \\cite{Bartolo2004}, \\cite{Renaux-Petel2009}, \\cite{Lewis2011}).  The most studied and constrained form of non-Gaussianity is the local ansatz parametrised by a factor $f_\\mathrm{NL}$, and predicted by several inflation models. Current constraints on local non-Gaussianity show that the CMB is consistent with being Gaussian at the 95$\\%$ C.L. \\citep{Komatsu2011}, and constraints on other shapes all show consistency with Gaussianity as well.  Given the current limits on primordial NG, astrophysical signals are the dominant contribution to non-Gaussianity. While Galactic emission and resolved sources may be accounted for by masking, unresolved sources and residuals have to be characterised. As opposed to primordial NG, radio sources NG has been detected, and was studied for the WMAP mission showing that it yields a non-zero flat angular bispectrum parametrised as $b_\\mathrm{src}$ or $b_\\mathrm{ps}$. The WMAP5 best estimate in the Q band is $b_\\mathrm{src} = 4.3 \\pm 1.3 \\mu K^3 \\mathrm{sr}^2$\\citep{Komatsu2009}. Characterising the NG signal from astrophysical components and more importantly from extra-galactic point sources is important for two main reasons: (i) to avoid mistaking it for a primordial contribution (and to allow the development of robust methods to isolate primordial NG) and (ii) to learn more about astrophysical processes, i.e. to go beyond the description of point sources by their number counts and their power spectrum.  The study of NG from extra-galactic point sources has been pioneered by \\cite{Argueso2003}, focusing mostly on the radio sources and including clustering. They showed that the point-source angular bispectrum is mostly flat at WMAP frequencies and dominates the CMB bispectrum in most configurations. \\cite{Babich2008} investigated the bias on the $f_\\mathrm{NL}$ estimator induced by radio sources, considering the modulation of their number density with ISW, of their magnification with lensing and of the flux cut-off with selection effects. \\cite{Serra2008} studied the bias on $f_\\mathrm{NL}$ due to radio sources, SZ-lensing and ISW-lensing bispectra, investigating the dependence of this bias with the resolution scale. Finally \\cite{Munshi2009} defined skew-spectra for cross-correlation analysis, derived estimators for the skew-spectra and applied it to secondary anisotropies.  In this paper we study the non-Gaussianity produced by infrared and radio point sources in the frequency range of the CMB from 30 to 350 GHz, based on numerical simulations by \\cite{Sehgal2010}. We investigate the frequency and configuration dependence of the angular bispectrum. We particularly focus on the non-Gaussianity from IR sources and their clustering term. We restrict the study to the simplest case of full sky maps without masks. Furthermore, we do not take into account noise and beam effects. Statistical isotropy of all fields considered will be assumed throughout this article. The case of noisy, partially masked maps will be addressed in future studies.  In section \\ref{sect:3pNGe} we review the estimator of the (binned) angular bispectrum and $f_\\mathrm{NL}$ and develop a parametrisation of the bispectrum to display and visualise it efficiently. In section \\ref{sect:prescription}, we develop a prescription to infer the bispectrum from the power spectrum for clustered sources and for different populations. In section \\ref{sect:resultsehgal} we use publicly available full-sky simulations of radio and infrared sources to compute and characterise their bispectrum at CMB frequencies and we compare them to the predictions from the prescription. We examine the configuration dependence of the point-source bispectra and study the bias they induce on the estimation of the primordial local non-Gaussianity in section \\ref{sect:ngcsqce}. We finally conclude and discuss our results in section \\ref{sect:concl}.   %************************************************************************** ",
        "conclusions": "\\label{sect:concl}  We have studied the non-Gaussianity produced by point sources in the frequency range of the CMB from 30 to 350 GHz. We have developed a simple and accurate prescription to infer the angular bispectrum from the power spectrum of point sources, considering different independent populations of sources, with or without clustering.  Using publicly available all-sky simulations of radio and IR sources, we have computed the full-sky binned bispectra for these two populations of sources. We have compared the measured bispectra to those predicted from our prescription and found a very good agreement between the two. We have displayed the angular bispectrum using a new parametrisation which highlights efficiently the configuration dependence.  We have characterised the angular bispectrum of the IR and radio sources showing the configuration dependence and the frequency behaviour. In particular and for the first time, we showed that the IR bispectrum peaks in the squeezed triangles and that the clustering of IR sources enhances the bispectrum values by several orders of magnitude on large angular scales $\\ell \\sim 100$.  The bispectrum of IR sources starts to dominate that of radio sources on large angular scales at 150 GHz, and it dominates the whole multipole range at 350 GHz.  Finally to illustrate the contamination of local CMB non-Gaussianity by point sources, we derive the bias on $f_\\mathrm{NL}$ induced by radio and IR sources, for WMAP or Planck-like angular resolutions. Radio sources produce a positive bias which is significantly reduced ($\\Delta f_\\mathrm{NL} < 1$ for $\\nu \\geq$ 150 GHz) by masking the sources above a given flux limit taken as the ERCSC cut. The form of the IR bispectrum mimics a primordial `local' bispectrum $f_\\mathrm{NL}$ on large angular scales. The IR sources produce a negative bias which becomes important for Planck-like resolution and at high frequencies ($\\Delta f_\\mathrm{NL} \\sim -6$ at 277 GHz and $\\Delta f_\\mathrm{NL} \\sim -$60-70 at 350 GHz). Most of the signal is associated with the clustering of faint IR sources. Therefore, the bias $\\Delta f_\\mathrm{NL}^\\mathrm{IR}$ is not reduced by masking sources above a flux limit but, in some cases, even increased due to the reduction of the shot-noise term.  Our analysis highlights the sensitivity of the bias on $f_\\mathrm{NL}$ to the experimental properties (maximum resolution and frequency range), the point-source models (clustering or no, resolved or not) and their scale dependence with respect to the CMB. For high resolution high frequency CMB experiments, primordial NG estimations need to take special care of astrophysical contaminations. One solution would be to estimate the primordial and astrophysical non-Gaussianity simultaneously.  %********************************************************************"
    },
    "1107/1107.4331_arXiv.txt": {
        "abstract": "{The supernova remnant Cassiopeia A is a prime candidate for accelerating cosmic ray protons and ions. Gamma rays have been observed at GeV and TeV energies, which indicates hadronic interactions, but they could also be caused by inverse-Compton scattering of low-energy photons by accelerated electrons.} {We seek to predict the flux of nuclear de-excitation lines from Cas A through lower-energy cosmic rays and to compare it with COMPTEL measurements.} {Assuming a hadronic origin of the high-energy emission, we extrapolate the cosmic ray spectrum down to energies of $10\\,\\mathrm{MeV}$, taking into account an equilibrium power-law momentum spectrum with a constant slope. We then calculate the nuclear line spectrum of Cassiopeia A, considering the most prominent chemical elements in the MeV band and their abundances as determined by X-ray spectroscopy.} {We show that the predicted line spectrum is close to the level of the COMPTEL sensitivity and agrees with conservative upper limits.} {} ",
        "introduction": "Cassiopeia A (Cas A) is one of the youngest known Galactic supernova remnants (SNR). Extensive observations in the radio, infrared, optical, and X-ray bands lead to an estimated explosion date of around 1680 AD \\citep{Thorstensen2001, Fesen2006}. Being the brightest radio source in our Galaxy, it was first discovered in 1948 \\citep{Ryle1948}. Its distance can be estimated to $3.4_{-0.1}^{+0.3}\\, \\mathrm{kpc}$ based on the combination of Doppler shifts and proper motions, whereas the angular size of $2.5'$ corresponds to a physical size of $2.34\\, \\mathrm{pc}$ \\citep{Baars1977}. The optical spectrum obtained from distant light echos of the original blast indicates that Cas A was a type IIb supernova of a $\\sim 15\\,M_\\odot$ main-sequence star and originated from the collapse of the helium core of a red supergiant that had lost most of its hydrogen envelope before exploding \\citep{Krause2008}.  The blast wave drives an initially outgoing forward shock that can be seen as a thin X-ray edge expanding at roughly $5000\\, \\mathrm{km\\, s^{-1}}$. The interaction with circumstellar material and the interstellar medium results in a reverse shock that is driven back into the outgoing ejecta and expanding at roughly half the rate of the forward shock \\citep{Gotthelf2001, DeLaney2003}. Emission at most wavelengths is dominated by a $30''$ thick ``bright ring'' where the ejecta are heated and ionized when they encounter Cas A's reverse shock. Viewed in X-rays, this shell consists of undiluted ejecta rich in O, Si, S, and Fe \\citep{Willingale2002, Laming2003, Lazendic2006}. Faint X-ray filaments outside of the shell mark the location of the forward shock where nonthermal X-ray synchroton radiation is produced by shock-accelerated electrons \\citep{Gotthelf2001, Vink2003}. Detected by HEGRA \\citep{Aharonian2001}, MAGIC \\citep{Albert2007} and VERITAS \\citep{Humensky2008}, Cas A was the first SNR verified in TeV gamma rays. Recent observations with Fermi-LAT in the GeV range do not rule out either a leptonic or a hadronic emission scenario: A combination of nonthermal bremsstrahlung and inverse-Compton emission on the one hand as well as neutral pion decays on the other hand could be responsible for the measured spectrum \\citep{Abdo2010}.  It has long been suggested that the shock waves associated with supernovae could be sites of acceleration of cosmic-ray particles \\citep{Baade1934}. The combination of diffusive shock acceleration (DSA) processes in SNR shocks with transport effects in our galaxy can in theory produce the observed power-law spectrum of cosmic rays up to the ``knee'' of about $10^{15}\\, \\mathrm{eV}$ \\citep{Krymskii1977, Axford1977, Bell1978, Bell1978a, Blandford1978, Drury1983, Jones1991}. To be consistent with the energy density of Galactic cosmic rays, the transfer of kinetic energy released in supernova explosions to cosmic rays must take place with an efficiency of $\\sim\\!10\\,\\%$ \\citep{Ginzburg1969}. According to the DSA model, the charged particles scatter off irregularities in the magnetic field and increase their momentum by a fraction of the shock velocity $v/c$ each round trip from the downstream to the upstream region and back again. Because of the feedback of the accelerated particles on the spatial profile of the flow velocity and therefore on the particle distribution itself, the DSA process can be significantly nonlinear. High-energy particles with considerable diffusion lengths sample a broader portion of the flow velocity profile and therefore experience a greater change in compression ratio than low-energy particles. That is why the particle spectrum and the resulting photon spectrum tend to flatten with increasing energy \\citep{Baring1999, Ellison2000, Bell2001, Ellison2005}. The correspondence of electron and proton spectra at high energies \\citep{Ellison2000} renders it extremely difficult to prove the acceleration of cosmic ray protons at SNR shocks directly. Though the presence of cosmic ray protons could be inferred by gamma rays resulting from collisions with ambient gas and subsequent pion decays, TeV gamma rays can also be produced by inverse-Compton scattering of cosmic microwave background photons with the accelerated cosmic-ray electrons in SNRs.  The detection of nuclear de-excitation gamma-ray lines could point a way out of this dilemma: The line spectra produced in energetic collisions of cosmic rays with the ambient medium provide an unambiguous proof for the hadronic scenario of cosmic ray acceleration in SNRs. In contrast to other proposed scenarios \\citep[cf.][]{Bozhokin1997}, we expect the excitation processes to happen near the proton acceleration site, so effects through Coulomb losses regarding the proton spectrum can be neglected here. The following sections are devoted to a short summary concerning the theoretical concepts of nuclear de-excitation processes as well as the derivation of the resulting gamma-ray fluxes in the MeV range with respect to current and future detection limits. Finally, we discuss the results and their implications in the last section. ",
        "conclusions": "We have shown that if cosmic rays are accelerated at the reverse shock in the Wolf-Rayet supernova remnant Cas A, as indicated by its strong GeV and TeV emission, the resulting flux of carbon and oxygen de-excitation lines is marginally detectable with COMPTEL. Precision spectroscopy with future MeV-missions such as GRIPS \\citep{Greiner2009} will put the theory of cosmic ray acceleration to a crucial test and will permit us to determine the overall efficiency of cosmic ray acceleration.  As a corollary, these measurements would also determine the yields of spallation products from nuclear collisions that affect the abundances of the cosmologically relevant light elements such as Li, Be, and B.  It must be noted that the prediction of the line flux rests on an extrapolation of the cosmic ray spectrum determined at high energies, which is in line with diffusive shock acceleration theory, but which is still subject to major systematic uncertainties associated to nonlinear effects such as shock broadening caused by back reactions."
    },
    "1107/1107.1452_arXiv.txt": {
        "abstract": "We present the results of fits of short-baseline neutrino oscillation data in 3+1 and 3+2 neutrino mixing schemes. In spite of the presence of a tension in the interpretation of the data, 3+1 neutrino mixing is attractive for its simplicity and for the natural correspondence of one new entity (a sterile neutrino) with a new effect (short-baseline oscillations). The allowed regions in the oscillation parameter space can be tested in near-future experiments. In the framework of 3+2 neutrino mixing there is less tension in the interpretation of the data, at the price of introducing a second sterile neutrino. Moreover, the improvement of the parameter goodness of fit is mainly a statistical effect due to an increase of the number of parameters. The CP violation in short-baseline experiments allowed in 3+2 neutrino mixing can explain the positive $\\bar\\nu_{\\mu}\\to\\bar\\nu_{e}$ signal and the negative $\\nu_{\\mu}\\to\\nu_{e}$ measurement in the MiniBooNE experiment. For the CP-violating phase we obtained two minima of the marginal $\\chi^2$ close to the two values where CP-violation is maximal. ",
        "introduction": "The recent agreement of MiniBooNE antineutrino data \\cite{1007.1150} with the short-baseline $\\bar\\nu_{\\mu}\\to\\bar\\nu_{e}$ oscillation signal observed several years ago in the LSND experiment \\cite{hep-ex/0104049} has opened an intense theoretical and experimental activity aimed at the clarification of the explanation of these observations in a framework compatible with the data of other neutrino oscillation experiments. Several short-baseline neutrino oscillation experiments did not observe neutrino oscillations and their data constraint the interpretation of the LSND and MiniBooNE antineutrino signal. However, there are other positive indications of short-baseline neutrino oscillations that may be taken into account: the reactor antineutrino anomaly \\cite{1101.2755}, in favor of a small short-baseline disappearance of $\\bar\\nu_{e}$, the Gallium neutrino anomaly \\cite{hep-ph/9411414,Laveder:2007zz,hep-ph/0610352,0707.4593,0711.4222,0902.1992,1005.4599,1006.2103,1006.3244,1101.2755}, in favor of a short-baseline disappearance of $\\nu_{e}$, and the MiniBooNE low-energy anomaly \\cite{0707.4593,0902.1992,1005.4599,1101.2755}. In this paper we consider only the reactor antineutrino anomaly, by taking into account the new calculation of reactor antineutrino fluxes in Ref.~\\cite{1101.2663}. We leave the discussion of the effects of the more controversial Gallium anomaly and MiniBooNE low-energy anomaly to a following article \\cite{Giunti-Laveder-IP-11}.  The results of solar, atmospheric and long-baseline neutrino oscillation experiments led us to the current standard three-neutrino mixing paradigm, in which the three active neutrinos $\\nu_{e}$, $\\nu_{\\mu}$, $\\nu_{\\tau}$ are superpositions of three massive neutrinos $\\nu_1$, $\\nu_2$, $\\nu_3$ with respective masses $m_1$, $m_2$, $m_3$. The measured solar (SOL) and atmospheric (ATM) squared-mass differences can be interpreted as \\begin{align} \\null & \\null \\Delta{m}^2_{\\text{SOL}} = \\Delta{m}^2_{21} = (7.6 \\pm 0.2) \\times 10^{-5} \\, \\text{eV}^2 \\quad \\text{\\protect\\cite{1010.0118}} \\,, \\label{SOL} \\\\ \\null & \\null \\Delta{m}^2_{\\text{ATM}} = |\\Delta{m}^2_{31}| = 2.32 {}^{+0.12}_{-0.08} \\times 10^{-3} \\, \\text{eV}^2 \\quad \\text{\\protect\\cite{1103.0340}} \\,, \\label{ATM} \\end{align} with $\\Delta{m}^2_{kj}=m_k^2-m_j^2$.  The completeness of the three-neutrino mixing paradigm has been challenged by the LSND \\cite{hep-ex/0104049} and MiniBooNE \\cite{1007.1150} observations of short-baseline $\\bar\\nu_{\\mu}\\to\\bar\\nu_{e}$ transitions at different values of distance ($L$) and energy ($E$), but approximately at the same $L/E$. Since the distance and energy dependences of neutrino oscillations occur through this ratio, the agreement of the MiniBooNE and LSND signals raised interest in the possibility of existence of one or more squared-mass differences larger than about 0.5 eV, which is much bigger than the values of $\\Delta{m}^2_{\\text{SOL}}$ and $\\Delta{m}^2_{\\text{ATM}}$. Hence, we are lead to the extension of three-neutrino mixing with the introduction of one or more sterile neutrinos which do not have weak interactions and do not contribute to the invisible width of the $Z$ boson \\cite{hep-ex/0509008}. In this paper we consider the simplest possibilities: 3+1 mixing with one sterile neutrino and 3+2 mixing with two sterile neutrinos.  The existence of sterile neutrinos which have been thermalized in the early Universe is compatible with Big-Bang Nucleosynthesis data \\cite{astro-ph/0408033,1001.4440}, with the indication however that schemes with more than one sterile neutrino are disfavored \\cite{1103.1261}, and cosmological measurements of the Cosmic Microwave Background and Large-Scale Structures if the neutrino masses are limited below about 1 eV \\cite{1006.5276,1102.4774,1104.0704,1104.2333,1106.5052}. Therefore, in this paper we consider squared-mass differences smaller than $10 \\, \\text{eV}^2$.  \\begin{figure}[t!] \\null \\hfill \\begin{minipage}[l]{0.2\\linewidth} \\begin{center} \\includegraphics*[bb=236 471 342 778, width=0.9\\linewidth]{fig-01a.eps} \\\\%[0.3cm] \\text{{\"normal\"}} \\end{center} \\end{minipage} \\hfill \\begin{minipage}[l]{0.2\\linewidth} \\begin{center} \\includegraphics*[bb=236 471 342 778, width=0.9\\linewidth]{fig-01b.eps} \\\\%[0.3cm] \\text{{\"inverted\"}} \\end{center} \\end{minipage} \\hfill \\null \\caption{ \\label{3+1} Schematic description of the two possible 3+1 schemes that we are considering, taking into account that $|\\Delta{m}^2_{21}| \\ll |\\Delta{m}^2_{31}| \\ll |\\Delta{m}^2_{41}|$. } \\end{figure} ",
        "conclusions": "In this paper we presented the results of fits of short-baseline neutrino oscillation data in 3+1 and 3+2 neutrino mixing schemes.  In the framework of 3+1 neutrino mixing schemes in Fig.~\\ref{3+1}, we confirm the strong tension between LSND and MiniBooNE antineutrino data and disappearance, KARMEN, NOMAD and MiniBooNE neutrino data discussed recently in Refs.~\\cite{0705.0107,1007.4171,1012.0267,1103.4570}. Since however the minimum value of the global $\\chi^2$ is rather good, one may choose to consider as possible 3+1 neutrino mixing, which can partially explain the data, taking into account its simplicity and the natural correspondence of one new entity (a sterile neutrino) with a new effect (short-baseline oscillations). Following this approach, we presented the results of the global fit in 3+1 neutrino mixing, which leads to the determination of restricted allowed regions in the mixing parameter space which can be explored in future $\\boss{\\nu}{\\mu}\\to\\boss{\\nu}{e}$ \\cite{CRUBBIA2011,1007.3228,1105.4984} $\\boss{\\nu}{e}$ disappearance \\cite{0907.3145,0907.5487,1006.2103,Porta:2010zz,1011.4509,1103.5307,CRUBBIA2011,Pallavicini2011,1105.4984} and $\\boss{\\nu}{\\mu}$ disappearance \\cite{CRUBBIA2011,1106.5685} experiments.  We also presented a global fit in the framework of the 3+2 neutrino mixing schemes in Fig.~\\ref{3+2}. We have shown that the tension between LSND and MiniBooNE antineutrino data and disappearance, KARMEN, NOMAD and MiniBooNE neutrino data is reduced with respect to the 3+1 fit, but is not eliminated (see Fig.~\\ref{3p2-sup-34}). Moreover, the improvement of the parameter goodness of fit with respect to that obtained in the 3+1 fit is mainly due to the increase of the number of oscillation parameters, as one can see from Tab.~\\ref{bef}. Hence it seems mainly a statistical effect.  The results of our 3+2 fit are compatible with those presented recently in Ref.~\\cite{1103.4570}, but we obtain a different indication for the best fit (see Tab.~\\ref{bef}). For the CP-violating phase we obtained two minima of the marginal $\\chi^2$ close to the two values where CP-violation is maximal.  In conclusion, we think that our results are useful for the discussion of the interpretation of the current experimental indications in favor of short-baseline neutrino oscillations and for the study of new experiments aimed at a clarification of the validity of these indications."
    },
    "1107/1107.1178_arXiv.txt": {
        "abstract": "For the very first time, we report the high frequency analysis of Cyg X-1 up to hard X-ray using SPI on-board {\\it INTEGRAL}. After analyzing the possible contribution from the background, and using {\\it INTEGRAL} archive from March 2005 to May 2008, Power Density Spectra (PDS) were obtained until 130 keV. First, we show that their overall shape is very similar to that observed at lower energies, as they are well described by sets of Lorentzians. The strength of this fast variability (up to 40 Hz) does not drop at high energy since we show that it remains at $\\sim 25\\%\\ \\rm rms$, even in the highest energy bands. Second, the hard X-ray variability patterns of Cyg X-1 are state dependent: the softer the spectrum (or the lower the hardness ratio), the lower the total fractional variability and the higher the typical frequencies observed. The strength of total variability as a function of energy and state is then investigated. By comparison with simultaneous and published {\\it RXTE}/PCA data, we showed that in the hard state, it remains quite constant in the 2-130 keV energy range. In our softer state, it is also flat until 50 keV and may increase at higher energy. The implications of this behavior on the models are then discussed. ",
        "introduction": "Black Hole Binaries (BHB), either transient (such as GX 339-4) or persistent (Cyg X-1), are very variable on long time scales (i.e., from weeks to months), in terms of flux, but also spectrally speaking: since the first spectra obtained (i.e., \\citealp{tananbaum1972} for Cyg X-1 and \\citealp{coe1976} for A0620-00), transitions between two main states were identified. The soft state, which spectrum is peaking in the soft X-ray band (i.e., $\\sim 1\\ \\rm keV$), is believed to be dominated by the emission of the optically thick accretion disk. For that latter reason, the soft state is also called ``Thermally Dominated'' \\citep{mcclintock2006}. In the hard state, the Spectral Energy Distribution (SED) is on the contrary peaking at few tens of keV, and is thought to be emitted by an optically thin component, the so-called ``corona''. Its origin, geometry, together with the heating processes of this component still remain unclear.    The observed radio behavior of those objects is also dependent on the spectral state \\citep{homan2005}: the canonical hard state is usually associated with the presence of a compact or sometimes extended steady jet (see e.g. \\citealp{stirling2001} for Cyg X-1), whereas the observation of relativistic blobs motion is associated with the intermediate state. In the soft state, on the contrary, the radio emission is quenched \\citep{fender1999}.  In order to explain such transitions, a paradigm has been set by \\cite{esin1997}, where in the soft state the disc is extending down to the vicinity of the black hole, whereas it is recessed in the hard state at hundreds of gravitational radii. The inner part is in turn replaced by an advection dominated flow. Observationally, the question to know if the disc is actually receding in the hard state is still under debate (see e.g. \\citealp{miller2006,reis2010}, but also \\citealp{gierlinski2008,cabanac2009,tomsick2009,done2010} ).  BHBs also exhibit strong signatures in the time domain at high frequencies. Their X-ray lightcurves are indeed characterized by erratic behavior and when analyzed in the Fourier domain, their Power Density Spectra (hereafter PDS) can show broad band aperiodic noise and/or broad incoherent peaks, the so called Quasi-Periodic Oscillations (QPO). The strength of those components, together with their typical frequency evolution, are closely dependent on the state \\citep{homan2005}. For a complete review of the X-ray binaries timing characteristics, see \\cite{vklis2004}.  Since its discovery in 1964 by an Aerobee rocket, and thanks to its persistent brightness due to the wind fed accretion process, Cyg X-1 has been extensively studied. The aforementioned spectral states are indeed well identified: \\cite{mcconnell2002} e.g., showed that the hard state (which is observed at a lower flux in the soft X-ray band), is well characterized by the power law component cutting off at $\\sim 100\\ \\rm keV$, whereas the soft state (observed at higher flux in the low energy bands) is dominated by the disk emission. In addition, in both soft and hard states, a high energy tail is also observed, which is usually interpreted as the emission of non thermal particles (see e.g., \\citealp{malzac2009} for a full SED modeling in presence of magnetic field, and the explanation of source state transition in terms of particle energy distribution variations).  However, as compared to the other known X-ray binaries, Cyg X-1 is in a  number of aspect quite different. No relativistic motion has been observed in the intermediate state \\citep{rushton2011} and the radio emission is not quenched in the soft state \\citep{rushton2011, zdziarski2011}. Its soft state is peculiar and sometimes called intermediate \\citep{belloni1996}. These differences may rely on the high mass nature of Cyg X-1 system, where wind fed accretion takes place, as compared to the vast majority of other BHBs which host low mass companions.  Cyg X-1 fast variability features were well studied by e.g. \\cite{pottschmidt2003}, making extensive use of the huge {\\it RXTE} data archive. They confirmed that the PDS could be well fitted by sets of Lorentzians, and they showed that the relative contributions of those sub-components were evolving with the states. The lower frequency Lorentzian is indeed suppressed relative to the second and third Lorentzian during the state transitions. Its energy behavior, or rms-spectrum, is rather constant in the hard state at $\\sim 25-40\\rm \\%\\ rms$.   There is however an important question which remains to address: are these high frequency variability features still present in the hard X-rays? At higher energies, a strong LFQPO (Low Frequency QPO) was indeed detected by SIGMA on board GRANAT \\citep{vikhlinin1994} in March 1990 and March 1991. Its typical peak frequency (between 0.04 and 0.07 Hz) was found to be similar to the broad band noise break frequency that was also detected in the power spectrum. However its behavior was very variable as the QPO, together with the broad band noise had disappeared one year later.  The energy dependency of the variability was also extensively studied in other sources: \\cite{rodriguez2004d, rodriguez2004c} found that a cut-off near 25 keV was probably present in GRS 1915+105 and XTE J1550-564 for the LFQPO rms-spectrum. Concerning the long time scales variability (from days to weeks), \\cite{zdziarski2002} used the {\\it RXTE}/ASM and CGRO/BATSE archives to study the variability in Cyg X-1 from 1.2 to 300 keV. They showed that this very low frequency variability was also quite high (i.e. 25\\% rms) and state dependent. Examining Cyg X-1 fast variability as a function of state, \\cite{gierlinski2010} (hereafter GZD10) showed that the rms spectrum was almost flat for the hard state whereas it is increasing with energy in the soft state. In intermediate states, it is decreasing with energy. Similar trend is observed in other BHBs (see e.g. \\citealp{gierlinski2005}). A tight and strong correlation was also found between the Comptonization amplification factor ($\\ell_{\\rm h}/\\ell_{\\rm s}$) and the frequency of the first Lorentzian $f_1$ as $\\ell_{\\rm h}/\\ell_{\\rm s}\\propto f_1^{-3/2}$.    The models to explain those variability features are still under debate, principally due to the various patterns observed and their evolution with the state. Its origin, i.e., is it coming from the disk, the corona, from a transition region or from the jet is even not well established. The most probable hypothesis is that it relies in an interaction between those components. Some successful attempts have however been built in order to reproduce some of the features observed: Concerning e.g. the LFQPO origin, the Accretion-Ejection Instability \\citep{tagger1999} was used  to explain the LFQPO frequencies and the disk geometry evolution \\citep{varniere2002}. \\cite{cabanac2010} built a model of an oscillating corona which addresses the QPO and the broad band noise typical frequencies correlation, the modeling of the PDS by radiative transfer simulation as well as the qualitative evolution of the frequencies with states : the softer the SED, the higher the frequencies. The observed increase of the absolute variability with the source flux, so called ``RMS-flux correlation'' \\citep{uttley2005}, was finally well modeled by \\cite{arevalo2006}, in the context of fluctuating-accretion model.  The purpose of this article is to study the fast timing behavior of Cyg X-1 at higher energies by using the high sensitivity of SPI (Spectrometer Onboard Integral, \\citealp{vedrenne2003,roques2003}), as its average effective area is $\\sim 500\\ \\rm cm^2$ in the 20 keV - 8 MeV band. Its timing ability is indeed unprecedented in those high energy bands, and it allowed, e.g., absolute timing of the Crab and determination of the hard X-rays to radio delay at  $\\sim 15 \\rm \\mu s$ resolution \\citep{molkov2010}. ",
        "conclusions": "SPI on-board {\\it INTEGRAL} was used successfully to study Cygnus X-1 variability at high energies. The high stability of the instrument allows to probe Cyg X-1 variability up to 100 Hz and down to 1 mHz. The first and main result of this study concerns the evolution of the variability at high energies: it is indeed demonstrated that Cyg X-1 is varying at a high level ($\\sim 25\\%$) up to 130 keV.  The shapes of the high energy PDS obtained are similar to that obtained at lower energies with other instruments as it is well described by subset of up to 4 Lorentzians. High frequency variability ($\\sim 40 \\rm Hz$) is observed in the 69-130 state averaged PDSs, showing for the first time that fast variability is present at high energy.   The parameters of the Lorentzians used are also evolving with the source state. For example, in order to explain the differences between the $HS_{SPI}$ and $SS_{SPI}$ frequencies found, the simplest solution is to associate those frequencies with given physical components, at certain typical radii in the disk or in the corona. In our case, the first Lorentzian is shifted from $1.6\\rm Hz$ in the $SS_{SPI}$ to $0.22\\rm Hz$ in the $HS_{SPI}$. If this component is linked to Keplerian motion, it would imply that the typical radius for this component increases by a factor $(0.22/1.6)^{-2/3}=3.8$ between the soft and hard state. In the framework where the variability is driven by waves propagating in the corona (see. e.g. \\citealp{cabanac2010}), the typical break frequency scales as the inverse of the external corona radius. Hence in that latter framework the corona external radius would be 7 times larger in our $HS_{SPI}$ than in the $SS_{SPI}$. It is interesting to note that the frequency shifts for $f_1$ follow the  difference between the hard and soft state fundamental planes found by \\cite{koerding2007}: this latter is based on the fundamental plane found by \\cite{mchardy2006} which links, for SS BHB and AGN, the mass accretion rate $\\dot{m}$, the first Lorentzian frequency $f_1$ and the black hole mass $M$ with $f_1 \\times M \\propto \\dot{m}$. Moreover, \\cite{koerding2007} have shown that this relation still holds in the HS, at price of changing the proportionality factor.  Based on an extensive analysis, they calculated that this factor was equal to 0.3 for HS BHB, whereas it was 2.2 in the SS. In our case, we note that the ratio between $\\Delta f_{1,\\ soft}$ and $\\Delta f_{1,\\ hard}$ is equal to 7, which is therefore very close to the proportionality factor ratio in both fundamental planes. % Such different tracks for HS and SS states in the fundamental plane is interpreted by \\citeauthor{koerding2007} as a sign that not only the accretion rate is driving the accretion process evolution. An additional parameter must be invoked to render this two different tracks in the Hard and Soft states.  The rms evolution with the hardness of the source that we observed at high energy is similar to the rms-hardness correlation seen in many XRB (see e.g., the Hardness-RMS Diagrams in \\citealp{belloni2010}). It reflects the known relation that the source gets more variable when it gets harder (see however some high variable soft state in GZD10, Fig. 2b). Our results support the fact that this relation still holds at high energies, giving strict constrains on emission models that intend to reproduce the observed variability.  The rms level we obtained ($\\sim 25\\%$) together with the frequency of the first Lorentzian ($\\Delta f_1\\sim0.2\\ \\rm Hz$) in our hard state  is compatible with the majority of the hard states found at lower energies with the PCA in GZD10 (see e.g. Figs 2b and 7a of that paper). If in GZD10 the hardest states rms reaches 35-40\\%, our hard state level, PDS and typical frequencies are compatible with bit softer ones. In particular, our hard state PDS (Fig.  \\ref{fig:psd-hard-soft}) is close in shape to the PCA PDS for observation 60090-01-19-00 (Fig 6b of GZD10), for which a rms spectrum is computed: the variability is rather constant around 25-30\\% rms on the 2-30 keV band, and adding our two high energy point would show that the rms remains at roughly the same level in the 2-130 keV band for this hard state. This rather constant rms spectrum in the hard state is emphasized by our simultaneous PCA/SPI analyse of the source, where the source was hard. If a break in the hard X-ray ($\\sim 80\\ \\rm keV$) is possible in this rms spectrum, its significance is rather marginal. It is also possible that the low duration of this observation and hence the lack of statistics does not allow to constrain a putative high frequency tail, which could account for extra variability. Independent results from \\cite{torii2011} obtained with {\\it Suzaku} showed that the hard state of Cyg X-1 rms is indeed rather constant in the 12-200 keV energy range.  Our $SS_{SPI}$ can be similarly associated with PCA observation id 70104-01-03-00 analysed by GZD10. The rms spectrum is in turn also quite constant between 2-35 keV, though its level is lower than in the previous hard state. Our SPI level of variability in the 27-49 keV band is consistent with the one obtained in the highest energy band of this PCA observation (15-20\\%) and it increases in the 69-130 keV (20-30\\%). An interesting interpretation of the rms spectra has been attempted by \\cite{gierlinski2005} and more recently in GZD10, with the following principles: any lightcurve variability measured by its fractional rms in a given energy band may be due to the variability of at least one of the parameters $p_i$ of the average spectrum. \\cite{gierlinski2005} and GZD10  used the combination of a multicolor disk ({\\verb diskbb }\\ \\citealp{mitsuda1984}), the {\\verb eqpair }\\  hybrid comptonization model \\citep{coppi1999} and a reflection component \\citep{magdziarz1995}, in order to model the average spectrum $F$. The effect on the output rms spectrum of each parameter variation is then investigated. Qualitatively, if a parameter $p_i$ only impacts on the normalization of a model, the spectrum will be shifted by a certain factor, independently from the energy. Hence, the variability of this parameter will lead to a flat rms spectrum. On the contrary, if another parameter $p_j$ (e.g. the photon index $\\Gamma$) changes the slope of the spectrum around a pivot, the factor by which the spectrum is shifted will not be energy independent: the further the working energy band is from the pivot, the higher the rms.  For the hard state, a flat rms spectrum is observed for XTE J1550-564 and XTE J1650-500 \\citep{gierlinski2005} and Cyg X-1 (GZD10). In those papers, it is proposed that the flat rms spectrum in the hard state may be reproduced by the number of soft photon variability or fluctuation of the accretion rate. Given that our rms spectrum behaviour remains flat until 130 keV in the hard state, we demonstrate that this interpretation still holds at high energies.  For our soft state examined here, the rms spectrum which is flat until 50 keV, may then increase at high energies. Following \\citep{gierlinski2005}, variability of the injection Lorentz factor $\\Gamma_{inj}$ may reproduce such rms spectrum increase at high energies, but a maximum should also be observed around 5-10 keV.  The above explanation of rms spectra however first depends on the underlying average model chosen, and second on the fact that only one parameter variation is considered. It is hence possible that some of the parameters are linked by underlying physical processes, such as e.g. the luminosity of the hard and soft components when $\\dot{m}$ is varying. As an example, \\cite{cabanac2010}, in the framework of an oscillating corona model, examined the impact of an adiabatic oscillation on the rms spectra. Both the corona temperature and the optical depth are then varying, which in turns leads to a Compton $y$ parameter variation. They predicted the presence of a minimum in the rms spectrum, due to the pivoting around the average spectrum. The position of the pivot depends on the relative peak between the soft photon source and the comptonized component.  We finally note that the high energy band power spectra that we obtained exhibits peaked-noise components, one of wich (the lower frequency one) is consistent with the QPO detected with SIGMA by \\cite{vikhlinin1994} in terms of peak frequency value. Those components are sharper in the higher than the lower energy PDSs (comparing e.g. Figs. \\ref{fig:psd-301-684} and \\ref{fig:psd-301-684-69-130}). This effect is also observed at lower energies with {\\it RXTE} (see e.g. Fig. 2 of \\citealp{cui1997}). If confirmed, it may therefore indicate that the Lorentzians  quality factor increases with energy. In other words, the coherence of the underlying variability processes would increase with energy, giving constraint on the emission models."
    },
    "1107/1107.4041_arXiv.txt": {
        "abstract": "{Models have been made of stars of a given mass which produce planetary nebulae which usually begin on the AGB (although they may begin earlier) and run to the white dwarf stage. While these models cover the so-called dredge-up phases when nuclear reactions occur and the newly formed products are brought to the surface, it is important to compare the abundances predicted by the models with the abundances actually observed in PNe.} {The aim of the paper is to determine the abundances in a group of PNe with uniform morphological and kinematic properties. The PNe we discuss are circular with rather low-temperature central stars and are rather far from the galactic plane. We discuss the effect these abundances have on determining the evolution of the central stars of these PNe.} {The mid-infrared spectra of the planetary nebulae \\object{NGC\\,1535}, \\object{NGC\\,6629}, \\object{He2-108}, and \\object{Tc1} (\\object{IC\\,1266}) taken with the {\\em Spitzer Space Telescope} are presented. These spectra were combined with the ultraviolet {\\em IUE} spectra and with the spectra in the visual wavelength region to obtain complete, extinction-corrected spectra.  The chemical composition of these nebulae is then found by directly calculating and adding individual ion abundances. For two of these PNe, we attempted to reproduce the observed spectrum by making a model nebula. This proved impossible for one of the nebulae and the reason for this is discussed. The resulting abundances are more accurate than earlier studies for several reasons, the most important is that inclusion of the far infrared spectra increases the number of observed ions and makes it possible to include the nebular temperature gradient in the abundance calculations.} {The abundances of the above four PNe have been determined and compared to the abundances found in five other PNe with similar properties studied earlier. These abundances are further compared with values predicted by the models of Karakas (2003). From this comparison we conclude that the central stars of these PNe originally had a low mass, probably between 1M\\smallsun~ and 2.5M\\smallsun. A further comparison is made with the stellar evolution models on the HR diagram, from which we conclude that the core mass of these PNe is between 0.56M\\smallsun~and 0.63M\\smallsun.} {A consistent picture of the evolution of this group of PNe is found that agrees with the predictions of the models concerning the present nebular abundances, the individual masses, and luminosities of these PNe. The distance of these PNe can be determined as well.} ",
        "introduction": "The evolution of planetary nebulae (PNe) has been studied for about four decades with the result that a general picture of the evolution is understood, but the details are still being debated. This is both because many of the parameters involved are poorly known and because the atmosphere of the central star is difficult to describe, including the effective temperature, radius, and mass loss rate. The (in many cases) uncertain distance enhances these difficulties. Furthermore, the interaction of the central star with the nebula is often unclear. This manifests itself in a difficulty explaining the ionization structure in some PNe.  The abundances of some elements change in the course of evolution of the central star. We discuss here primarily the elements helium, nitrogen, carbon, and oxygen, although neon, argon, sulfur, chlorine, and sometimes other elements are determined as well. These elements are ejected by the star and are measured in the nebula. These abundances provide complementary information which aid in understanding the evolution and which can eventually confirm or limit the results of models of the evolution.  We are concerned here with abundances found in PNe whose central stars are rather bright and whose morphology is usually described as round or elliptical. We present new abundances for four such nebulae: NGC1535, NGC6629, Tc1 (IC\\,1266), and He2-108. The abundances in these nebulae will be combined with the results found earlier in five other PNe, which are also excited by bright stars and which have similar (round or elliptical) morphology. These PNe have similar spatial properties as well: they are at rather high galactic latitudes. We can therefore speak of a class of nebulae. This class has been extensively studied earlier but from a different point of view. The central stars of these PNe are among the brightest known in the visible. It is therefore possible to observe these spectra with very high resolution. The profiles of hydrogen and helium lines can then be analyzed with the initial goal of obtaining the effective temperature (T$_{eff}$) and gravity of these stars. The final goal is to use these quantities to study the mass and luminosity (thus the evolution) of these stars, but additional information is required to do this.  One of the early extensive studies of the line profiles is that of \\cite{mendez,mendez2}. These authors take high resolution optical spectra of 24 PNe and interpret these spectra using NLTE plane-parallel static model atmospheres which contain only hydrogen and helium. The values of T$_{eff}$ and the stellar helium abundance are determined by fitting the observed HeI and HeII line profiles while the gravity (log g) is found by fitting the hydrogen line profiles (usually H$\\gamma$ but sometimes H$\\beta$). Then, making use of theoretical evolution diagrams which plot T$_{eff}$ against log g for different values of the stellar mass M$_{s}$/M\\smallsun, ~the value of M$_{s}$/M\\smallsun~is determined. Using this mass and the value of gravity, the angular radius of the exciting star is determined. Combining this value with the previously found temperature and the observed magnitude corrected for extinction, the distance to the nebulae is found. In the course of time this group has improved the model atmospheres used by first including the sphericity of the atmosphere \\citep{kudritzki1} and then by including other elements besides H and He in the atmosphere \\citep{kudritzki2}. These improved models do not change the stellar masses or PNe distances substantially. Both the masses and the distances determined in this way, however, are suspect. Quoting \\cite{kudritzki2} `the masses determined seem systematically larger than white dwarf masses and some of the objects ... have unrealistically high masses'. Concerning the distances: for the 18 cases where the distance thus found may be compared with statistical distance \\citep{stang2,cahn} they are always higher, in 11 cases more than 60\\% higher.  Recently \\cite{pauldrach} used a different approach. By including models losing mass they are able to use 'observed' mass loss rates and terminal velocities to obtain the stellar masses and distances. These masses and distances are even higher than found by \\cite{kudritzki2} and are shown to be implausibly high by \\cite{napiwotzki}. Two arguments are used. Because high mass central stars are probably descended from high mass progenitors they should have the properties of high mass stars. First they should belong to the disk population of the galaxy which is characterized by small scale heights perpendicular to the galactic disk. Secondly such high mass star should have produced and dredged up substantial amounts of both nitrogen and helium. Napiwotzki concludes that on both of these points the high masses found by \\cite{pauldrach} are unlikely.  This suggests the following step: the abundances of those elements produced by the various stars should be accurately determined. These abundances may then be used to determine the mass of the star in question by comparing the observed abundance with the abundances determined by nucleosynthesis model calculations made in the course of evolution of stars of different masses. For these models, the calculations made by \\cite{karakas} are used. These masses can then be compared with those found by \\cite{kudritzki2}. Because the masses used in the calculations of Karakas are the initial stellar masses, and the masses given by \\cite{kudritzki2} are the present (almost final) stellar masses, a ratio between the initial and final mass must be known. For this secondary information the values given by \\cite{weide} will be used. Unfortunately the uncertainties present in each of these steps will accumulate so that definite conclusions are difficult to draw. It is however important that central star masses can be derived from accurate nebular abundances.  The main purpose of this paper is to obtain accurate abundances for four of the nebulae discussed by \\cite{kudritzki}.  The most important reason that this can be achieved is the inclusion of the mid-infrared spectrum taken with the IRS spectrograph of the {\\em Spitzer Space Telescope} (Werner et al. 2004). The reasons for this have been discussed in earlier papers \\citep[e.g. see][]{pott1,pott2, pott4, bernard}, and can be summarized as follows: 1) The intensity of the infrared lines is not very sensitive to the electron temperature nor to possible extinction effects.  2) Use of the infrared line intensities enable a more accurate determination of the electron temperature for use with the visual and ultraviolet lines. 3) The number of observed ionization stages is doubled.  For two of the PNe a second method of determining the abundances is attempted using a nebular model. This has several advantages. First it provides a physical basis for the electron temperature determination. Secondly it permits abundance determination for elements which are observed in only one, or a limited number of ionic stages. This is true of Mg, Fe, P and Cl, which would be less reliably determined without a model. A further advantage of modeling is that it provides information on the central star and other properties of the nebula.  A disadvantage of modeling is that there are more unknowns than observations and some assumptions must be made especially concerning the geometry and the form of the radiation field of the exciting star. The difficulties we encountered in determining a model for NGC\\,1535 are discussed below.  Abundance determination for these nebulae have been made earlier and will be compared with our values in the individual sections.  This paper is structured as follows. In Sect.\\, 2, 3, 4 and 5 the observations of NGC\\,1535, NGC\\,6629, Tc1 and He2-108 will be discussed and the resultant abundances will be given and compared to earlier results. In Sect.\\,7, the models for NGC\\,1535 and Tc1 are presented and discussed.  In Sect.\\,8 the evolutionary state of the nebulae is discussed.  Finally, our conclusions are given in Sect.\\,9. ",
        "conclusions": "With the help of Spitzer infrared spectra the abundances in four PNe have been determined. These PNe are all excited by rather low temperature central stars and have similar morphological and kinematic properties: they are all nearly round and are rather far from the galactic plane. We are able to show that these nebulae have rather similar abundances of helium, oxygen, nitrogen, carbon and other elements. The resultant abundances are summarized in the first four lines of Table 17. We then show that five other PNe with low temperature central whose abundances have been determined using Spitzer infrared spectra and have the same or similar morphological and kinematic properties also have the same or similar abundances.  By comparing these abundances with those predicted by nucleosynthesis models by \\cite{karakas} it is deduced that these PNe originate from stars of initial mass between 1M\\smallsun~and about 2.5M\\smallsun, which according to \\cite{weide} correspond to core masses of between 0.56M\\smallsun~and 0.63M\\smallsun. These values of core masses are compared with those determined from stellar evolution theory using the observed temperature of the central star and the measured age of the nebula. The core masses thus found are consistent with each other. Two details reinforce this consistency. First, the higher core masses from the evolutionary theory are found in PNe which have high C/O abundance ratios as the models of \\cite{karakas} predict. Secondly the distances found from the stellar evolution are in general values expected from statistical distance scales.  There are a number of uncertainties which must still be considered. First, it is not understood why the central star temperature measured in NGC\\,2392 is so much lower than that found from the nebula. Second, distances found by some researchers are different than found here. For example, the distances given by \\cite{kudritzki2} for 6 of the PNe listed are 50\\% higher than we have found. This could be due to errors in their determination of the stellar gravity from uncertain line profiles. Furthermore the expansion distances found for two of the nebulae (NGC\\,3242 and IC\\,2448) are at least a factor of 2 lower than we have found here. This should be carefully considered in the future."
    },
    "1107/1107.2384_arXiv.txt": {
        "abstract": "We derive a method to reconstruct Gaussian signals from linear measurements with Gaussian noise. This new algorithm is intended for applications in astrophysics and other sciences. The starting point of our considerations is the principle of minimum Gibbs free energy which was previously used to derive a signal reconstruction algorithm handling uncertainties in the signal covariance. We extend this algorithm to simultaneously uncertain noise and signal covariances using the same principles in the derivation. The resulting equations are general enough to be applied in many different contexts. We demonstrate the performance of the algorithm by applying it to specific example situations and compare it to algorithms not allowing for uncertainties in the noise covariance. The results show that the method we suggest performs very well under a variety of circumstances and is indeed qualitatively superior to the other methods in cases where uncertainty in the noise covariance is present. ",
        "introduction": "The problem of signal inference consists of reconstructing a set of parameters or even a continuous field $s$ from some data set $d$, which is influenced in some way by the signal, \\begin{equation} \\label{eq:datamodel-general} d=f(s)+n. \\end{equation} Two problems will arise. First, the function $f$ may not be invertible and, second, the noise term $n$ will not be known. In the Bayesian framework, one uses prior information on the signal and the noise term to calculate a best estimate for the true signal realization or, ideally, the whole probability distribution for the signal given the prior information and the information contained in the data.  Symmetry considerations and knowledge about the underlying physics of the signal and the measurement process may restrict the class of priors that one has to consider. They might, however, still contain some free parameters that then become part of the inference problem. The case in which the signal covariance contains uncertain parameters was tackled in \\cite{ensslin_frommert-2011}, producing a whole class of filters for this problem. The filter that we extend in this work was reproduced in \\cite{ensslin_weig-2010}, where the principle of minimum Gibbs free energy was introduced (cf. also Sect.~\\ref{sec:derivation}), and successfully applied in an astrophysical setting in \\cite{oppermann-2011a}.  Here, we focus on the case where we can assume zero-mean Gaussian priors both for the signal and for the noise. The priors are therefore completely characterized by the respective covariance matrices. Our goal is to extend the study of \\cite{ensslin_weig-2010} to the case in which both the signal covariance and the noise covariance contain parameters that are not known a priori. This is motivated mainly by applications from the field of astrophysics. The theory and resulting filter formulas, however, are of general applicability. Gaussian noise, e.g., is omnipresent in nearly every area of the natural sciences and the situation in which its variance is not precisely known should be a rather common one.  Previous work dealing with the problem of unknown noise variance has mainly dealt with specific applications. One of these applications is the field of image reconstruction. Here, it is usually assumed that the measured picture is the sum of the underlying signal and a white Gaussian noise term. Often, it is further assumed that the noise level, i.e. its variance, is the same in every image pixel. A comparison of different algorithms for noise estimation under these assumptions can be found e.g. in \\cite{olsen-1993}. An example for an algorithm allowing for inhomogeneous noise is presented in \\cite{starck-1998}, where a wavelet transform of the image is applied and the lack of correlated noise is exploited. Most of these algorithms, however, are not derived by rigorous statistical calculations but rather by a combination of intuition and experience.  From a mathematical viewpoint, the problem of an unknown noise prior has received some attention in the theory of density deconvolution, which deals with the inference of the probability density for a signal from measurements with additive noise. Here, the signal is usually assumed to consist of independent identically distributed variables. The case of Gaussian noise with unknown variance has been considered e.g. by \\cite{koltchinskii-2000} and \\cite{schwarz-2010}.  In this work, we create a general setting with well defined assumptions and a traceable derivation of a general filter formula within a Bayesian framework, not loosing sight of its applicability. Our result can accomodate a host of different assumptions and models, such as correlated or uncorrelated noise. It allows for a distinction between the data space and the signal space, with possibly different numbers of degrees of freedom.  The remainder of the paper is organized as follows. In Sect.~\\ref{sec:notation} we introduce our model for the measurement process and the notation that is required. The derivation of the filter formulas follows in Sect.~\\ref{sec:derivation}. We then demonstrate the usefulness of our filter by applying it in a set of mock observational situations in Sect.~\\ref{sec:application} and discuss the implications in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}   Using the formalism of minimum Gibbs free energy we have extended the critical filter algorithm, developed in \\cite{ensslin_frommert-2011} and \\cite{ensslin_weig-2010}, to an algorithm that allows for uncertainties both in the signal covariance and in the noise covariance. We have demonstrated the performance of our algorithm, Eqs.~\\eqref{eq:gcf}, by applying it to a set of mock observations on the sphere, as well as in a simple one-dimensional test case.  These applications have shown that the extended critical filter performs outstandingly if only a few individual data points have a misestimated error bar. However, even in a case where large portions of the data are affected, the algorithm was shown to perform inarguably better than the critical filter, using a fixed -- and faulty -- assumption about the noise statistics. We have also compared the results to those obtained from a Wiener filter reconstruction, using the correct power spectrum, which is known to be optimal if the assumptions about the noise statistics are correct. This filter was demonstrated, however, to lead to reconstructed maps that are much further from the true signal than the results of the extended critical filter in some cases where the assumptions are not correct.  The choice of the two-sphere as the space on which our signal is defined was motivated by astrophysical applications, where we could think of the signal as an all-sky field or some quantity defined on the surface of a star or a planet. Applications in other fields of physics are abundant. However, it should be noted that there is nothing special about the sphere. We could equally well have chosen a more-dimensional euclidean space, using the power spectrum defined in Fourier space instead of the angular power spectrum, as we have done in the one-dimensional scenario.  Furthermore, our choice of the identity operator as response matrix was made only on account of simplicity. It allowed us to represent the data in the same fashion as the signal. It should be clear, however, that the derived filter formulas, Eqs.~\\eqref{eq:gcf}, are valid for any response matrix, even a singular one. Applications of the critical filter with non-trivial response matrices were presented in \\cite{ensslin_frommert-2011} and \\cite{oppermann-2011a} and such a response would not pose a problem for the extended version of the filter.  The problem of signal reconstruction with some uncertainty in the noise variance is certainly one of general interest. There are several ways in which uncertainty in the noise variance might arise. It may be due to questionable assumptions that enter in the calculation of the error bars of the data. Another possibility is that it arises from the definition of the signal itself. The quantity of interest may only be part of what has been measured in the first place in which case the rest of the data would be noise with essentially unknown variance. All these factors come together in the reconstruction problem considered in \\cite{oppermann-2011a}. An extension of that reconstruction, using additional sets of data and the improved algorithm presented here, is planned \\citep{oppermann-2011c}.  It should be noted, however, that even with the extended critical filter, some knowledge about either the parameters in the signal covariance or the ones in the noise covariance is needed to arrive at a sensible reconstruction. Leaving them both completely free would lead to a degeneracy between signal and noise that cannot be resolved. Only by assigning an informative prior to at least one of the two sets of parameters is this degeneracy broken. Furthermore, the functional bases in which the signal and noise covariances are diagonal, i.e. their eingenspaces, need to differ to allow for a separation of noise from signal."
    },
    "1107/1107.5047_arXiv.txt": {
        "abstract": "Measurements of cosmic microwave background (CMB) anisotropies constrain isocurvature fluctuations between photons and nonrelativistic particles to be subdominant to adiabatic fluctuations. Perturbations in the relative number densities of baryons and dark matter, however, are surprisingly poorly constrained. In fact, baryon-density perturbations of fairly large amplitude may exist if they are compensated by dark-matter perturbations, so that the total density remains unchanged. These compensated isocurvature perturbations (CIPs) leave no imprint on the CMB at observable scales, at linear order. B modes in the CMB polarization are generated at reionization through the modulation of the optical depth by CIPs, but this induced polarization is small.  The strongest known constraint $\\lesssim 10\\%$ to the CIP amplitude comes from galaxy-cluster baryon fractions. Here it is shown that modulation of the baryon density by CIPs at and before the decoupling of Thomson scattering at $z\\sim 1100$ gives rise to CMB effects several orders of magnitude larger than those considered before.  Polarization B modes are induced, as are correlations between temperature/polarization spherical-harmonic coefficients of different $lm$. It is shown that the CIP field at the surface of last scatter can be measured with these off-diagonal correlations.  The sensitivity of ongoing and future experiments to these fluctuations is estimated. Data from the WMAP, ACT, SPT, and Spider experiments will be sensitive to fluctuations with amplitude $\\sim5-10\\%$. The Planck satellite and Polarbear experiment will be sensitive to fluctuations with amplitude $\\sim3\\%$.  SPTPol, ACTPol, and future space-based polarization methods will probe amplitudes as low as $0.4\\%-0.6\\%$. In the cosmic-variance limit, the smallest CIPs that could be detected with the CMB are of amplitude $\\sim0.05\\%$. ",
        "introduction": "\\label{iref}  The concordance cosmological model posits a nearly scale-invariant spectrum of primordial density fluctuations with adiabatic initial conditions, for which the ratios of neutrino, photon, baryon, and dark-matter number densities are homogeneous.  The simplest inflationary models predict adiabatic fluctuations \\cite{Guth:1982ec,Linde:1982uu,bardeen_iso,Hawking:1982cz,mukhanov_iso,starobinsky_iso}, and adiabatic fluctuations are consistent with measurements of cosmic microwave background (CMB) temperature/polarization anisotropies \\cite{wmap7_komatsu_a} and the clustering of galaxies \\cite{sdss_lrg,sdss_pspec}.  Isocurvature perturbations are fluctuations in the ratios of number densities of various particle species.  They are produced in topological-defect models for structure formation \\cite{branden_topo} and in more complicated models of inflation \\cite{linde_iso,linde_mukhanov,mukhanov_b,axenides,langlois_riazuelo,langlois_double_iso}. CMB temperature anisotropies limit the amplitude of baryon isocurvature perturbations (fluctuations in the baryon-to-photon ratio) \\cite{peebles_iso_a,peebles_iso_b} and CDM isocurvature perturbations (fluctuations in the dark-matter--to--photon ratio) \\cite{burns_axion,seckel_turner,hu_iso_a,hu_iso_b} to be $\\lesssim13\\%$ of the total perturbation amplitude \\cite{boomerang_iso_limits,planck_iso,wmap7_komatsu_a,wmap7_komatsu_b,tommaso,zunckel_iso,bucher_dunkley_a,kawasaki_iso,beltran_viel_iso,seljak_iso,riazuelo_iso_constraint,takahashi_iso_nongauss}.  Our intuition thus suggests the matter in the early Universe was very smoothly distributed.  It therefore comes as somewhat of a surprise to learn that perturbations in the baryon density can be almost arbitrarily large, as long they are compensated by dark-matter perturbations such that the total nonrelativistic matter density remains unchanged \\cite{gordon_pritchard,holder}. These compensated isocurvature perturbations (CIPs) thus obey \\begin{eqnarray} \\rho_{\\rm c}\\delta_{\\rm c}^{\\rm CI}+\\rho_{\\rm b}\\delta^{\\rm CI}_{\\rm b}=0,~~~\\delta_{\\gamma}^{\\rm CI}=0,\\end{eqnarray} where $\\delta_{\\rm c}$, $\\delta_{\\rm b}$, and $\\delta_{\\rm \\gamma}$ are fractional energy density perturbations in the dark matter, baryons, and photons, respectively, while $\\rho_{\\rm c}$ and $\\rho_{\\rm b}$ are the homogeneous dark matter and baryon densities.  CIPs induce no curvature perturbation at early times, and they therefore leave the photon density---and thus large-angle CMB fluctuations---unchanged at linear order.  CIPs induce baryon motion through baryon-pressure gradients, but these motions occur only at the baryon sound speed which, at the time when Thomson scattering first decouples ($z\\sim 1100$, \\textit{decoupling} hereafter), is $(v/c)\\sim (T/m_p)^{1/2}\\sim (\\mathrm{eV}/\\mathrm{GeV})^{1/2}\\sim 10^{-4.5}$.  The effects of these motions on CMB temperature and polarization anisotropies thus occur only on distances smaller than $\\sim10^{-4.5}$ times the sound horizon at decoupling or CMB multipole moments $l\\sim 10^6$ \\cite{gordon_pritchard,gordon_lewis_curvaton,challinor_compensate}, scales far smaller than those probed by CMB experiments.  The effect of CIPs on galaxy surveys is also believed to be small \\cite{gordon_pritchard}.  Big-bang nucleosynthesis (BBN) and galaxy-cluster baryon fractions constrain the CIP perturbation amplitude to be $\\lesssim10\\%$ \\cite{holder}. Measurements of fluctuations in 21-cm radiation from atomic hydrogen during the dark ages may be sensitive to these perturbations \\cite{barkana,challinor_compensate,gordon_pritchard,sekigu}, but these measurements are a long way in the future.  In Ref.~\\cite{holder}, it was shown that, although CIPs produce no observable effect on the CMB at linear order in perturbation theory, they modulate the CMB fluctuations produced by adiabatic perturbations.  In particular, it was shown that B modes in the CMB polarization are produced by the angular modulation in the reionization optical depth induced by the CIP.  Here, we consider the additional effects on CMB fluctuations that arise from modulation of the baryon density by CIPs at and before decoupling.  CIPs modulate the free-electron density. They thus change the photon diffusion length and thickness of the surface of last scattering (SLS) on different patches of sky. CIPs also change the weight of the baryon-photon plasma and thus the details of the acoustic-peak structure in the CMB power spectrum.  Variation in the baryon density from one region on the sky to another thus leads to a modulation of the small-scale power spectrum from one region of sky to another. This induces B modes in the polarization and nontrivial higher-order correlations in the temperature/polarization map analogous to those induced by variations of other cosmological parameters \\cite{sigurdson_alpha} and those induced by weak gravitational lensing \\cite{lewis_lens_review}.  As we show below, the effects of CIPs on CMB fluctuations from decoupling are several orders of magnitude larger than those from reionization, and so the CMB should provide a far more sensitive probe of CIPs than envisioned in Ref.~\\cite{holder}.  We therefore follow through and develop the formalism required to look for CIPs with the CMB.  To do so, we write down the minimum-variance estimators that can be constructed from a CMB temperature-polarization map for the CIP field $\\Delta(\\hat n)$ as a function of position $\\hat n$ on the sky.  We evaluate the noise with which the CIP field can be reconstructed and estimate the signal-to-noise with which a scale-invariant spectrum of CIPs may be detected with various experiments.  We conclude that data from WMAP, Spider, ACT, and SPT are sensitive to CIP amplitudes of $\\sim 5-10\\%$. The Planck satellite \\cite{planck_bluebook} and Polarbear experiment are sensitive to CIP amplitudes as small as $\\sim 3\\%$. Upcoming ground-based polarization experiments (ACTPol \\cite{actpol} and SPTPol \\cite{dvorkin_smith_b,sptpol}) or a post-Planck CMB-polarization experiment along the lines of the proposed EPIC experiment \\cite{cmbpol} could detect fluctuations of $\\sim 0.4\\%-0.6\\%$. In the cosmic-variance limit, sensitivity to fluctuations of amplitude $\\sim 0.05\\%$ is possible.  Our principal motivation in studying CIPs is curiosity: can we determine empirically, rather simply assume, that the primordial baryon fraction is homogeneous and traces the dark matter? Still, there may be theoretical motivation as well.  For example, curvaton models for inflation may generate CIPs \\cite{lyth_ungarelli_wands,gupta_malik_wands,gordon_lewis_curvaton,enqvist_subdom_curvaton}, with amplitudes approaching the regime detectable by EPIC \\cite{gordon_pritchard}. It may also be that recent models \\cite{Kaplan:2009ag,buckley_hylogenesis,cladogenesis,wimp_baryo,heckman_dm,barymorphosis} that connect the baryon asymmetry and dark-matter density have implications for CIPs. Additionally, the techniques introduced in this paper could be used to empirically disentangle a CDM isocurvature fluctuation from a baryon isocurvature fluctuation, using CMB data. These modes are usually treated as degenerate in the analysis of CMB observations.  In Ref. \\cite{grin_prl}, we presented our basic conclusions. Here we present in detail our results, their derivation, and the computational methods used. We calculate the induced temperature anisotropies in Sec.~\\ref{tempsec} and the induced polarization anisotropies in Sec.~\\ref{polaniso}. In Sec.~\\ref{pspec_predict}, we compute the expected corrections to CMB power spectra for a scale-invariant spectrum of CIPs and compare the B-mode power spectrum induced by CIPs at decoupling with that induced at reionization. In Sec.~\\ref{qestsecbig}, we construct minimum-variance estimators for CIPs.  We then assess in Sec.~\\ref{exprop} the sensitivity of ongoing and upcoming experiments to CIPs, and we conclude in Sec.~\\ref{conc}.  Useful relations involving tensor spherical harmonics are presented in Appendix \\ref{tensor_on_sphere}. Numerical derivatives of transfer functions are discussed in Appendix \\ref{deriv_pspec_sec}. Second-order harmonic expansions for CMB transfer functions are derived in Appendix \\ref{secdercorr}.  Throughout, we use as our fiducial cosmological parameters those from Ref.~\\cite{wmap7_komatsu_a}. ",
        "conclusions": "\\label{conc}  Compensated isocurvature perturbations provide an intriguing empirical possibility for large-amplitude departures from homogeneity in the early Universe.  The current constraint, $\\lesssim 10\\%$, to the amplitude of such perturbations is surprisingly weak.  While Ref.~\\cite{holder} has pointed out that there may be CMB signatures induced by CIPs at reionization, we have shown here that the CMB effects of CIPs at the surface of last scatter would be several orders of magnitude larger.  We then calculated the full two-point temperature/polarization correlations induced by CIPs on the CMB and developed the minimum-variance estimators for measuring the CIP field with the CMB.  The WMAP satellite may be sensitive to a scale-invariant spectrum of CIPs, but only if the CIP amplitude is close to its current upper bound. In the future, sensitivity to CIP amplitudes as small as $\\sim 3 \\%$ may be achieved by instruments in operation and $\\sim 0.1\\%$-level fluctuations accessible in the near future with precise ground- and space-based polarization experiments that are under construction (ACTPol and SPTPol) or in conceptual development (EPIC).  Many steps must be taken before such measurements can be implemented with real data.  Techniques to deal with partial-sky coverage and realistic instrumental noise properties must be developed, but these techniques should be similar to those being developed already to measure the effects of weak gravitational lensing on the CMB.  Likewise, techniques must be developed to distinguish the off-diagonal correlations induced by CIPs from those induced by weak gravitational lensing (e.g., \\cite{meng}), which should be comparable in amplitude if the CIP field $\\Delta$ is comparable to the lensing potential $\\phi$, i.e, $\\sim 1\\%$.  Although lensing is already included in our reconstruction noise estimates [$\\sigma_{\\Delta_{L}}$ in Eq. (\\ref{toterr})], it might also induce bias in the reconstruction of the CIPs. Gravitational lensing of the CMB will induce correlations of similar form to Eqs.~(\\ref{correlaa}) and (\\ref{correlab}), with the distinction that the functions $S_{l l^{\\prime}}^{L,{\\rm X X^{\\prime}}}$ will be different for lensing than for CIPs. In the case of lensing, these functions describe the remapping of CMB observables on a lensing-deflected sky. In the case of CIPs, these functions describe the detailed physical dependence of CMB anisotropies on the baryon density. If Eq.~(\\ref{toterr}) is applied to the extension of Eqs.~(\\ref{correlaa})-(\\ref{correlab}) that includes gravitational lensing, a bias will be induced in the measurement of $\\Delta_{LM}$.  The differing forms of $S_{l l^{\\prime}}^{L,{\\rm X X^{\\prime}}}$ will allow lensing and CIPs to be disentangled. Using a straightforward generalization of the estimator in Eq.~(\\ref{toterr}) that includes terms for both CIPs and lensing, the bias induced by lensing on CIP measurements may be removed, simultaneously reconstructing the lensing and CIP fields. This is analogous to the estimators discussed in Ref. \\cite{meng}, where it is shown that if reionization is patchy because of inhomogeneity in the distribution of ionized bubbles around the first sources, the contributions of patchy reionization and lensing may be distinguished. In future work, we will explicitly compute the bias in CIP searches that will be induced by lensing and will construct the estimators that disentangle lensing from CIPs. As in Ref. \\cite{meng}, we expect that the signal-to-noise of the biased and unbiased estimators should be nearly the same, and so gravitational lensing should not affect the signal-to-noise estimates of this paper.  The measurements we propose in this paper offer a precise test of how closely the primordial baryon and dark-matter distributions are matched and approach the CIP amplitudes allowed in curvaton models. Moreover, if future CMB experiments detect subdominant isocurvature fluctuations between matter and radiation, the techniques developed in this work could disentangle contributions from the baryon and cold dark-matter (CDM) isocurvature modes. Even greater gains in sensitivity are theoretically possible if the cosmic-variance limit is approached at high $l$ by future experiments. We are optimistic that, in the near future, we will learn just how well baryons trace dark matter in the early Universe."
    },
    "1107/1107.1933_arXiv.txt": {
        "abstract": "As part of the NASA-CNES agreement, the NASA Star and Exoplanet Database (NStED) serves as the official US portal for the public CoRoT data products. NStED is a general purpose archive with the aim of providing support for NASA's planet finding and characterization goals. Consequently, the NASA Exoplanet Science Institute (NExScI) developed, and NStED adapted, a periodogram service for CoRoT data to determine periods of variability phenomena and create phased photometric light curves. Through the NStED periodogram interface, the user may choose three different period detection algorithms to use on any photometric time series product, or even upload and analyze their own data. Additionally, the NStED periodogram is remotely accessed by the CoRoT archive as part of its interface. NStED is available at {\\bf http://nsted.ipac.caltech.edu}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.3088_arXiv.txt": {
        "abstract": "For many decades the determination of accurate fundamental parameters for stars (masses, radii, temperatures, luminosities, etc.) has mostly been the domain of eclipsing binary systems.  That has begun to change as long-baseline interferometric techniques have improved significantly, and powerful new instruments have come online.  This paper will review the status of the field, and in particular how the knowledge of precise stellar properties helps us understand stars. Main-sequence stars similar to the Sun are by far the best studied, but much remains to be done for other kinds of objects such as early-type as well as late-type stars including brown dwarfs, evolved stars, metal-poor stars, and pre-main sequence stars.  Progress is illustrated with several examples of how interferometry has contributed significantly in some of these areas. ",
        "introduction": "The determination of the fundamental properties of stars, i.e., their masses, radii, effective temperatures, luminosities, etc., is a classic discipline in Astronomy that goes back more than a century. To many, this area of work may not seem as appealing or fashionable as other topics such as cosmology and extrasolar planet research, which seem to garner most of the attention these days. It is worth keeping in mind, though, that there are few areas in modern Astrophysics that do not rely to some extent on our knowledge of the basic properties of stars, and this includes cosmology and extrasolar planets.  Perhaps the most important application of precise measurements of the mass and other stellar characteristics is to improve our understanding of stellar structure and stellar evolution, by comparing the observations against predictions from current models. Numerous examples of such comparisons may be found in the literature \\citep[see, e.g.,][]{Pols:97, Lastennet:02, Hillenbrand:04}, and have shown that while our knowledge of the physics of stars appears to be in reasonably good shape for solar-type stars, the agreement in other parts of the H-R diagram is not as good. More practical applications include the estimate of the total mass in stellar clusters making use of the well-known mass-luminosity relation determined empirically from observations of binary systems (the only ones allowing a dynamical measurement of the mass). Binary stars with well-determined properties serve also as valuable distance indicators, not only in the Milky Way but also in other galaxies (SMC, LMC, M31, M33), and have become an important tool for establishing the large-scale structure of the Universe.  In recent years the field of extrasolar planets has highlighted the importance of understanding stars. For transiting exoplanets, for example, in which the observables are the transit light curve and typically also a spectroscopic radial-velocity curve, the masses and radii of these objects cannot be determined independently of those of the parent star. Even if the inclination angle is precisely known in these cases, the value of the planet mass derived from the spectroscopic orbit scales as $M^{2/3}$, where $M$ is the mass of the host star. Similarly, the measurement of the depth of the transit events does not immediately yield the radius of the planet, but only the \\emph{ratio} between the planet radius and the stellar radius. It is essential to know the properties of stars in order to learn about planets. Thus, in this era of deep questions about the possibility of life on other worlds, and the frequency of planets like our own in the Universe, stellar Astronomy has become relevant again.  In this review I will outline the techniques applied to measure stars, and the status of the field of fundamental stellar parameter determination. I will also illustrate some of the contributions from long-baseline interferometry, which is the subject of this Workshop. ",
        "conclusions": ""
    },
    "1107/1107.3041_arXiv.txt": {
        "abstract": "Until now, axisymmetric, $\\alpha$-disc models have been adopted for calculations of the chemical composition of protoplanetary discs. While this approach is reasonable for many discs, it is not appropriate when self-gravity is important. In this case, spiral waves and shocks cause temperature and density variations that affect the chemistry. We have adopted a dynamical model of a solar-mass star surrounded by a massive (0.39 M$_{\\odot}$), self-gravitating disc, similar to those that may be found around Class 0 and early Class I protostars, in a study of disc chemistry.  We find that for each of a number of species, e.g.\\  H$_2$O, adsorption and desorption dominate the changes in the gas-phase fractional abundance; because the desorption rates are very sensitive to temperature, maps of the emissions from such species should reveal the locations of shocks of varying strengths. The gas-phase fractional abundances of some other species, e.g.\\ CS, are also affected by gas-phase reactions, particularly in warm shocked regions. We conclude that the dynamics of massive discs have a strong impact on how they appear when imaged in the emission lines of various molecular species. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5451_arXiv.txt": {
        "abstract": "We study the evolution of Super Star Cluster (SSC) winds driven by stellar winds and supernova (SN) explosions. Time-dependent rates at which mass and energy are deposited into the cluster volume, as well as the time-dependent chemical composition of the re-inserted gas, are obtained from the population synthesis code Starburst99. These results are used as input for a semi-analytic code which determines the hydrodynamic properties of the cluster wind as a function of cluster age. Two types of winds are detected in the calculations. For the quasi-adiabatic solution, all of the inserted gas leaves the cluster in the form of a stationary wind. For the bimodal solution, some of the inserted gas becomes thermally unstable and forms dense warm clumps which accumulate inside the cluster. We calculate the evolution of the wind velocity and energy flux and integrate the amount of accumulated mass for clusters of different mass, radius and initial metallicity. We consider also conditions with low heating efficiency of the re-inserted gas or mass loading of the hot thermalized plasma with the gas left over from star formation. We find that the bimodal regime and the related mass accumulation occur if at least one of the two conditions above is fulfilled. ",
        "introduction": "\\label{sec:intro}  Super star clusters are young compact objects observed in many starburst and interacting galaxies in a variety of wavelengths \\citep{ 1992AJ....103..691H, 1993AJ....106.1354W, 1995ApJ...446L...1O, 2005ApJ...619..270M, 2006MNRAS.370..513S, 2007ApJ...668..168G, 2008A&A...492....3G, 2011ApJ...729..111W}. With masses $10^5 - 10^7$~M$_\\odot$ and ages $\\lesssim 10^7$~yr they are expected to include large numbers of massive stars which lose substantial fractions of their mass via stellar winds and supernova explosions.   \\citet[hereafter CC85]{1985Natur.317...44C} studied the hydrodynamics of the gas re-inserted by massive stars into the cluster interior using an adiabatic spherically-symmetric model. They assumed that the mechanical energy of stellar winds and supernovae ejecta is thermalized in random collisions and the gas within the cluster is heated up to $\\sim 10^7$~K. The resulting high pressure drives the cluster wind for which CC85 found a stationary hydrodynamic solution. They assumed that the mass and the thermal energy are inserted uniformly at rates $\\dot{M}_\\mathrm{SC}$ and $L_\\mathrm{SC}$, respectively, into a sphere (cluster) of radius $R_\\mathrm{SC}$. They showed, that under such assumptions, a stationary wind can only be obtained if the flow velocity equals zero at the cluster center and reaches the sound speed exactly at the cluster border. Super star cluster winds were studied further using analytical and numerical models by many authors including \\citet{ 2000ApJ...536..896C, 2001ApJ...559L..33R,2003ApJ...590..791S,2006ApJ...643..186T}.   It was found by \\citet{2004ApJ...610..226S} that the adiabatic approximation becomes inadequate for very massive and compact clusters. The authors showed that the stationary solution of the cluster wind does not exist for clusters with $L_\\mathrm{SC}$ larger than a critical value $L_\\mathrm{crit}$. This is because the total energy input rate, $L_\\mathrm{SC}$, is proportional to the cluster stellar mass, $M_\\star$, while the energy losses from the hot gas due to radiation are proportional to $M_\\star^2$ (since cooling is proportional to the second power of the gas density which is proportional to $M_\\star$). \\citet{2004ApJ...610..226S} showed how $L_\\mathrm{crit}$ depends on the star cluster parameters and \\citet{2007A&A...471..579W} founded an approximate analytical formula for $L_\\mathrm{crit}$.  Clusters with $L_\\mathrm{SC} > L_\\mathrm{crit}$ were studied by means of 1D hydrodynamic simulations by \\citet{2007ApJ...658.1196T}, who showed that such clusters evolve in the bimodal hydrodynamic regime. In such a case, the cluster is divided by the stagnation radius, $R_\\mathrm{st}$, into two qualitatively different regions. The stationary wind solution still exists in the outer region $r > R_\\mathrm{st}$, with the wind velocity being zero at $R_\\mathrm{st}$ and reaching the sound speed at $R_\\mathrm{SC}$. In the region $r < R_\\mathrm{st}$, on the other hand, the thermal instability sets in and random parcels of gas cool down to $\\sim 10^4$~K (further cooling is prevented by the intense stellar radiation). Consequently, the warm regions are compressed into dense clumps by repressurizing shocks driven by the surrounding hot gas. Clusters in the bimodal regime were studied further by \\citet{2008ApJ...683..683W} who used 2D hydrodynamics to follow the clump formation, and to estimate the fraction of the re-inserted matter which leaves the cluster as a wind and the fraction which accumulates inside the stagnation radius and possibly leads to secondary star formation \\citep{2005ApJ...628L..13T}.   It was suggested that two-component supersonic recombination line profiles often detected in young and massive SSCs \\citep{2007ApJ...668..168G, 2008A&A...489..567B, 2007AJ....133..757H} and compact dense HII regions overlapping young SSCs \\citep{2006MNRAS.370..513S} may present the observational manifestation for such bimodal regime \\citep{2010ApJ...708.1621T, 2007ApJ...669..952S, 2009ApJ...700..931S}. In both cases the calculations require the shocked gas temperature to be lower than that predicted by the CC85 model as it is also the case when the model predicted diffuse X-ray emission is compared to the observed values \\citep{2003MNRAS.339..280S}. Two different processes which may decrease the intercluster gas temperature have been discussed in the literature: the efficiency with which the kinetic energy of stellar winds and SNe is thermalized, and the additional mass loading into the hot gas inside the cluster \\citep{2003MNRAS.339..280S, 2004A&A...424..817M, 2007A&A...471..579W, 2007ApJ...669..952S, 2009ApJ...700..931S, 2010ApJ...711...25S}. In this work we do not discuss details related to those two processes, however, we introduce two free parameters $\\eta_\\mathrm{he}$ and $\\eta_\\mathrm{ml}$ and show how the results depend on their values.  Previous works on clusters in the bimodal regime use the energy and mass deposition rates $L_\\mathrm{SC}$ and $\\dot{M}_\\mathrm{SC}$ as free parameters. In this work, we calculate time-dependent $L_\\mathrm{SC}(t)$ and $\\dot{M}_\\mathrm{SC}(t)$ using the stellar population synthesis code Starburst99 \\citep{1999ApJS..123....3L} for a cluster with a given stellar mass, $M_\\star$, and initial stellar metallicity, $Z_0$. Subsequently, we insert $L_\\mathrm{SC}(t)$ and $\\dot{M}_\\mathrm{SC}(t)$ into our semi-analytic code to determine the evolutionary properties of the cluster wind. We also calculate whether the cluster spends some time in the bimodal regime and estimate the amount of re-inserted gas which becomes thermally unstable and accumulates inside the cluster. The Starburst99 code also provides us with the time evolution of the re-inserted gas chemical composition. The chemical composition is an important parameter as the cooling rate depends on it. This work effectively replaces the three functions of time $L_\\mathrm{SC}(t)$, $\\dot{M}_\\mathrm{SC}(t)$ and $Z(t)$ (metallicity of the cluster wind), with the two constant parameters: mass of the star cluster $M_\\star$ and its initial metallicity $Z_0$.  The paper is organized as follows: in \\S\\ref{sec:model} we describe the semi-analytic code used for the calculation of the cluster wind and the way how it utilizes results of the Starburst99 code. In \\S\\ref{sec:results} we show results for a reference model with $M_\\star = 10^6$~M$_\\odot$ and $R_\\mathrm{SC} = 3$~pc (\\S\\ref{ssec:proto}) and give the dependence of results on the cluster mass, the cluster radius and the initial stellar metallicity (\\S\\ref{ssec:dep}). In \\S\\ref{sec:conclusions} we summarize our conclusions.   \\begin{figure} \\plotone{fig1.eps} \\caption{Energy (solid line, left y-axis) and mass (dashed line, right y-axis) deposition rates calculated by the Starburst99 code for the reference model $M_\\star = 10^6$~M$_\\odot$, $R_\\mathrm{SC} = 3$~pc, $Z_0 = $Z$_\\odot$, $\\eta_\\mathrm{he} = 1$ and $\\eta_\\mathrm{ml} = 0$. } \\label{fig:depos} \\end{figure}  \\begin{figure} \\plotone{fig2.eps} \\caption{The evolution of the metallicity of the hot shocked gas inside the cluster with $Z_0 = Z_\\odot$. The solid and dashed lines show the metallicity without ($\\eta_\\mathrm{ml} = 0$) and with ($\\eta_\\mathrm{ml} = 19$) mass loading, respectively.} \\label{fig:chem} \\end{figure} ",
        "conclusions": ""
    },
    "1107/1107.5721_arXiv.txt": {
        "abstract": "We present a novel approach of identifying the Milky Way (MW) and Andromeda (M31) progenitors that could be visible as LAEs at $z \\sim 6$: we couple a snapshot from the Constrained Local UniversE Simulations (\\texttt{CLUES}) project, that successfully reproduces the MW and M31 galaxies situated in their correct environment, to a Lyman Alpha Emitter (LAE) model. Exploring intergalactic medium (IGM) ionization states ranging from an almost neutral to a fully ionized one, we find that including (excluding) the effects of clustered sources the first local group progenitor appears as a LAE for a neutral hydrogen fraction $\\chi_{HI} =0.4$ ($\\chi_{HI}= 0.1$). This number increases to 5 progenitors each of the MW and M31 being visible as LAEs for $\\chi_{HI} = 10^{-5}$; the contribution from clustered sources is crucial in making many of the progenitors visible in the Ly$\\alpha$, for all the ionization states considered. The stellar mass of the local group LAEs ranges between $10^{7.2-8} {\\rm M_\\odot}$, the dust mass is between $10^{4.6-5.1} {\\rm M_\\odot}$ and the color excess $E(B-V)=0.03-0.048$. We find that the number density of these LAEs are higher than that of general field LAEs (observed in cosmological volumes) by about two (one) orders of magnitude for $\\chi_{HI}=10^{-5}$ (0.4). Detections of such high LAE number densities at $z \\sim 6$ would be a clear signature of an over-dense region that could evolve and resemble the local group volume at $z=0$. ",
        "introduction": "\\label{intro} The search for high-redshift LAEs has advanced tremendously since the days of their first successful detections, e.g. Lowenthal et al. (1991). Advances in instrument sensitivity, refined selection techniques and specific LAE spectral signatures have led to the confirmed detections of hundreds of LAEs in a wide redshift range, between $z \\sim 2.2-6.6$ (e.g. Cowie \\& Hu 1998; Steidel et al. 2000; Malhotra et al. 2005; Shimasaku et al. 2006; Kashikawa et al. 2006; Hu et al. 2010).  Over the past few years, LAEs have been used extensively as probes of reionization, high redshift galaxy evolution and the dust content of early galaxies (Dijkstra et al. 2007; Kobayashi et al. 2007; Dayal et al. 2009; Finkelstein et al. 2009; Dayal, Ferrara \\& Saro 2010; Dayal, Maselli \\& Ferrara 2011). However, scant effort has been devoted to establishing a link between such high-redshift observations ($z \\sim 6$) to those at $z = 0$, especially in the context of the local group.  The first such attempt was made by Salvadori, Dayal \\& Ferrara (2010), who coupled the semi-analytic code, {\\tt GAMETE} (Salvadori, Schneider \\& Ferrara 2007; Salvadori \\& Ferrara 2009), which successfully reproduces most of the observed MW and dwarf satellite properties at $z=0$, to a LAE model developed in Dayal et al. (2009) and Dayal, Ferrara \\& Saro (2010). Though a powerful statistical tool, the semi-analytic nature of {\\tt GAMETE} meant that the effects of clustered sources on the visibility of LAEs could not be properly considered; to alleviate such a problem to some extent, all the calculations presented therein were carried out assuming the IGM to be completely ionized.  In this work, we extend the calculations presented in Salvadori, Dayal \\& Ferrara (2010). We use a $z \\sim 6$ snapshot from the {\\tt CLUES} runs, which are state of the art constrained simulations that build up one of the possible merger histories of the local group, including the MW and M31 galaxies. This snapshot is coupled to a LAE model presented in Dayal, Ferrara \\& Saro (2010) and Dayal, Maselli \\& Ferrara (2011), that reproduces a number of data sets accumulated for high-z LAEs. The main advantage of using these simulations lies in the fact that they provide information on both the physical properties and the spatial positions of all the progenitors of the local group at any given redshift. This allows us to calculate the effects of clustered sources on the visibility of such progenitors in the Ly$\\alpha$, for IGM ionization states ranging from a completely neutral ($\\chi_{HI}=0.99$) to a fully ionized ($\\chi_{HI}=10^{-5}$) one.  The main questions we aim to answer with such an approach are: (a) how reionized does the volume containing the local group have to be at $z \\sim 6$ for any of its progenitors to be visible as a LAE?, (b) where do such LAEs lie with respect to the underlying dark matter (DM) distribution?, and (c) what are the physical properties of the progenitors that are visible as LAEs, and how do they compare to those of field LAEs observed in cosmological searches at $z \\sim 6$, e.g. by Shimasaku et al. (2006). ",
        "conclusions": "\\label{conc} We present the first work at high-z ($z \\sim 6$) that couples state of the art high resolution gas-dynamical simulations of the local group run within the \\texttt{CLUES} framework, with a LAE model to identify the local group progenitors (designated LAE$_{LG}$) that could be visible as LAEs at a time when the Universe was only about 1 Gyr old. The main results from this study are now summarized:  \\begin{itemize}  \\item The first LAE$_{LG}$, which is a progenitor of the MW, appears at an average IGM neutral hydrogen fraction of $\\chi_{HI}=0.4$ (0.1) including (neglecting) the effects of clustered sources. As $\\chi_{HI}$ decreases to $10^{-5}$, 5 progenitors each of the MW and M31 become visible as LAE$_{LG}$.  \\item $T_\\alpha$ shows a large scatter, ranging between $0.35-0.65$ at $\\chi_{HI}=10^{-5}$, due to the varying photoionization rate contributions from the clustered sources of the LAE$_{LG}$; clustering is imperative in making many progenitors visible in the Ly$\\alpha$ even for a fully ionized IGM.  \\item The LAE$_{LG}$ lie at the low mass end of the field LAEs: they have $M_* \\sim 10^{7.2-8} {\\rm M_\\odot}$, $M_{dust} \\sim 10^{4.6-5.1} {\\rm M_\\odot}$ and $E(B-V) \\sim 0.03-0.048$. On the other hand, field LAEs are much larger and have $M_* \\sim 10^{8-10.5} {\\rm M_\\odot}$, $M_{dust} \\sim 10^{4-7.2} {\\rm M_\\odot}$ and an average $E(B-V) =0.14$.  \\item Finally, our results suggest an observable imprint of such high-z LAE$_{LG}$: the number density of LAE$_{LG}$ is about two (one) orders of magnitude higher than that of field LAEs for $\\chi_{HI} = 10^{-5}$ ($\\chi_{HI}=0.4$). Such {\\it high number-densities} at $z \\sim 6$ would be excellent signatures of a region that could resemble the local group volume at $z \\sim 0$. However, the extent to which such a result depends on the varying simulation resolutions still needs to be explored in more detail.   \\end{itemize}  The main caveats involved in this study are with regards to the distribution of dust (homogeneous/clumpy) in the interstellar medium of these high-z progenitors (e.g. Finkelstein et al. 2009; Dayal, Ferrara \\& Saro 2010; Dayal, Maselli \\& Ferrara 2011), and the effects of peculiar velocities (inflows/outflows into/from the galaxies) that can have enormous impact on $T_\\alpha$ (Verhamme et al. 2006; Dayal, Maselli \\& Ferrara 2011). Uncertainties also remain on the value of the escape fraction of \\HI ionizing photons used. As discussed in the above mentioned works, it is unclear to what extent these effects can modify the observed Ly$\\alpha$ luminosity. It is hoped that a combination of the expected data from upcoming missions such as ALMA, MUSE and JWST will be fundamental in resolving such issues."
    },
    "1107/1107.2115_arXiv.txt": {
        "abstract": "{We derive a sufficient condition for realizing meta-stable de Sitter vacua with small positive cosmological constant within type IIB string theory flux compactifications with spontaneously broken supersymmetry. There are a number of `lamp post' constructions of de Sitter vacua in type IIB string theory and supergravity. We show that one of them -- the method of `K\\\"ahler uplifting' by F-terms from an interplay between non-perturbative effects and the leading $\\alpha'$-correction -- allows for a more general parametric understanding of the existence of de Sitter vacua. The result is a condition on the values of the flux induced superpotential and the topological data of the Calabi-Yau compactification, which guarantees the existence of a meta-stable de Sitter vacuum if met. Our analysis explicitly includes the stabilization of all moduli, i.e. the K\\\"ahler, dilaton and complex structure moduli, by the interplay of the leading perturbative and non-perturbative effects at parametrically large volume. } ",
        "introduction": "\\setlength{\\parskip}{0.2cm}  String theory is a candidate for a fundamental theory of nature, providing at the same time a UV-finite quantum theory of gravity and unification of all forces and fermionic matter. Mathematical consistency requires string theory to live in a ten dimensional space-time, and a description of our large four-dimensional physics thus necessitates compactification of the additional six dimensions of space.  The need for compactification confronts us with two formidable consequences: Firstly, even given the known internal consistency constraints of string theory, there are unimaginably large numbers of 6d manifolds available for compactification. Secondly, many compact manifolds allow for continuous deformations of their size and shape while preserving their defining properties (such as topology, vanishing curvature, etc) -- these are the moduli, massless scalar fields in 4d. This moduli problem is exacerbated if we wish to arrange for low-energy supersymmetry in string theory, as compactifications particularly suitable for this job -- Calabi-Yau manifolds -- tend to come with hundreds of complex structure and K\\\"ahler moduli.  Therefore, a very basic requirement for string theory to make contact with low-energy physics is moduli stabilization -- the process of rendering the moduli fields very massive. Moreover, as supersymmetry is very obviously broken -- and so far has not been detected -- ideally, moduli stabilization should tolerate or even generate supersymmetry breaking. And finally, the process should produce a so-called meta-stable de Sitter (dS) vacuum with tiny positive cosmological constant, so as to accommodate the observational evidence for the accelerated expansion of our universe by dark energy~\\cite{Perlmutter:1998np,Riess:1998cb,Komatsu:2010fb}.  The task of moduli stabilization and supersymmetry breaking has recently met with considerable progress, which is connected to the discovery of an enormous number~\\cite{Bousso:2000xa,Giddings:2001yu,Kachru:2003aw,Susskind:2003kw,Douglas:2003um} of stable and meta-stable 4d vacua in string theory. The advent of this \ufffdlandscape\ufffd~\\cite{Susskind:2003kw} of isolated, moduli stabilizing minima marks considerable progress in the formidable task of constructing realistic 4d string vacua.  There are several methods of moduli stabilization. The first one uses supersymmetric compactifications of string theory on a Calabi-Yau manifold, and the strong gauge dynamics of gaugino condensation in the `racetrack' mechanism to stabilize the dilaton and several of the bulk volume and complex structure moduli~\\cite{Derendinger:1985kk,Dine:1985rz,deCarlos:1992da}. Recently, this method has been applied to supersymmetric compactifications of M-theory on $G_2$-manifolds, where the structure of the manifolds allows for the racetrack superpotential to generically depend on all the moduli of the compactification~\\cite{Acharya:2006ia}.  The second, more recent, method relies on the use of quantized closed string background fluxes in a given string compactification. These flux compactifications can stabilize the dilaton and the complex structure moduli of type IIB string theory compactified on a Calabi-Yau orientifold supersymmetrically~\\cite{Giddings:2001yu}. The remaining volume moduli are then fixed supersymmetrically by non-perturbative effects, e.g. gaugino condensation on stacks of D7-branes~\\cite{Kachru:2003aw}. The full effective action of such fluxed type IIB compactifications on Calabi-Yau orientifolds was derived in~\\cite{Grimm:2004uq}. In type IIA string theory on a Calabi-Yau manifold all geometric moduli can be stabilized supersymmetrically by perturbative means using the larger set of fluxes available~\\cite{DeWolfe:2005uu}.  If the moduli are stabilized supersymmetrically, parametrically small and controlled supersymmetry breaking can happen, e.g, by means of inserting an anti-D3-brane into a warped throat of the Calabi-Yau~\\cite{Kachru:2003aw}, by D-terms originating in magnetic flux on a D7-brane~\\cite{Burgess:2003ic}, or dynamically generated F-terms of a matter sector~\\cite{Lebedev:2006qq}. This process is known as `uplifting' and allows for dS vacua with extremely small vacuum energy by means of fine-tuning the ${\\cal O}(100)$ independent background fluxes available in a typical Calabi-Yau compactification~\\cite{Bousso:2000xa,Kachru:2003aw}. The reliability of this last step of uplifting supersymmetric AdS vacua without unstabilized moduli into a dS vacuum is still under discussion. Some of the points in question e.g. concern the fact that the existence of D7-brane D-terms as well as F-terms from hidden matter sectors are very model dependent, rendering statistical sweeps over large sets of compactifications difficult. Supersymmetry breaking and uplifting by a warped-down anti-D3-brane also remains under ongoing discussion on whether its presence can be completely described in a probe approximation or causes dangerous non-normalizable perturbations to the compact geometry~\\cite{DeWolfe:2008zy,Bena:2009xk,Dymarsky:2011pm,Blaback:2011nz,Bena:2011wh,Burgess:2011rv}. Very recently, the use of internal $F_2$ gauge flux on a CY threefold in heterotic string theory has been used to stabilize all geometric  moduli except the dilaton and one K\\\"ahler modulus in a supersymmetric Minkowski vacuum~\\cite{Anderson:2010mh,Anderson:2011cz}.  Alternatively, in non-Calabi-Yau flux compactifications of type IIB or IIA string theory, all geometric moduli can be stabilized perturbatively in a non-supersymmetric way using a combination of background fluxes, D-branes, orientifold planes, and negative curvature. Examples here are flux compactifications of type IIB with 3-form fluxes on a product of Riemann surfaces~\\cite{Saltman:2004jh} and almost Calabi-Yau 4-folds in F-theory~\\cite{Dong:2010pm}, type II compactifications with generalized fluxes on manifolds of $SU(3)$ (see, e.g., the reviews~\\cite{Grana:2005jc,Douglas:2006es}), as well as of type IIA with fluxes on a product of two 3d nil manifolds~\\cite{Silverstein:2007ac}. The ingredients used typically lead to scalar potential dominated by three perturbative terms with alternating signs, which depend as varying power laws on the dilaton and the geometric moduli. Such a `3-term structure'  structure generically allows for tunable dS vacua~\\cite{Saltman:2004jh,Silverstein:2007ac}. Supersymmetry is generically broken in these perturbative mechanisms of moduli stabilization at a high scale, which typically is the Kaluza-Klein (KK)-scale. The geometric and flux part of these type IIA compactifications were studied in more detail in~\\cite{Caviezel:2008tf,Flauger:2008ad,Danielsson:2009ff,Douglas:2010rt,Wrase:2010ew,Danielsson:2010bc,Underwood:2010pm,Shiu:2011zt,Burgess:2011rv}. The conclusion there so far seems to be that in absence of the KK5-branes used in~\\cite{Silverstein:2007ac} (which play a similar role as 'explicit' supersymmetry breaking objects as the anti-D3-brane in~\\cite{Kachru:2003aw}) there are no stable dS vacua. A complete analysis including the effects of the KK5-branes in the language of~\\cite{Caviezel:2008tf,Flauger:2008ad,Danielsson:2009ff,Douglas:2010rt,Wrase:2010ew,Danielsson:2010bc,Underwood:2010pm,Shiu:2011zt,Burgess:2011rv} still remains open.  Finally, in type IIB flux compactifications on Calabi-Yau manifolds there are constructions of a `hybrid' type, where fluxes fix the complex structure moduli and the dilaton supersymmetrically, but the volume moduli are stabilized non-super-symmetrically by an interplay of non-perturbative effects on D7-brane stacks and the leading perturbative correction at ${\\cal O}(\\alpha'^3)$ in type IIB~\\cite{Becker:2002nn}, or by perturbative corrections to the K\\\"ahler potential alone.  Examples for the latter consist of the Large-Volume-Scenario (LVS)~\\cite{Balasubramanian:2005zx}, stabilization by perturbative corrections to the K\\\"ahler potential of the volume moduli alone~\\cite{Berg:2005yu,Berg:2007wt,Cicoli:2007xp,vonGersdorff:2005bf} which are uplifted by D7-brane D-terms~\\cite{Parameswaran:2006jh}, and the method of `K\\\"ahler uplifting'~\\cite{Westphal:2006tn,Balasubramanian:2004uy}.  For `K\\\"ahler uplifted' dS vacua, an interplay between the leading perturbative correction at ${\\cal O}(\\alpha'^3)$ and a non-perturbative effect in the superpotential serves to generate a dS vacuum with supersymmetry spontaneously broken by an F-term generated in the volume moduli sector. For some recent reviews on flux compactifications and the associated questions of the landscape of string vacua and string cosmology ensuing from the meta-stable dS vacua, with a much more complete list of references, please see~\\cite{Douglas:2006es,McAllister:2007bg,Baumann:2009ni}.  `K\\\"ahler uplifting' has the benefit of generating meta-stable dS vacua in terms of just background 3-form fluxes, D7-branes and the leading perturbative ${\\cal O}(\\alpha'^3)$-correction, data which are completely encoded in terms of the underlying F-theory compactification on a fluxed Calabi-Yau fourfold.   In addition, supersymmetry is spontaneously broken at a scale of order of the inverse Calabi-Yau volume, measured in string units this is typically $\\sim M_{\\rm GUT}$ here, and still below the KK-scale), by an F-term generated in the volume moduli sector. No extra anti-branes, D-terms or F-term generating matter fields are needed or involved. The existing analysis of these models consists of including manifestly the dilaton and one complex structure modulus~\\cite{Westphal:2006tn}.  Therefore, in this paper we develop a method towards a rigorous analytical understanding of `K\\\"ahler uplifting' driven by the leading ${\\cal O}(\\alpha'^3)$ correction to the K\\\"ahler potential of the volume moduli. Our derivation will be carried out in the presence of an arbitrary number $h^{2,1}$ of complex structure moduli. A large value of 3-cycles $h^{2,1}={\\cal O}(100)$  is a prerequisite to use the associated 3-form fluxes for the required fine-tuning of the cosmological constant.  Note the relationship between the supersymmetric KKLT-type AdS vacua~\\cite{Kachru:2003aw} (prior to uplifting) with the flux superpotential tuned small, the SUSY-breaking LVS-type AdS vacua~\\cite{Balasubramanian:2005zx} (again, prior to uplifting), and the SUSY-breaking `K\\\"ahler uplifted' AdS/dS vacua~\\cite{Westphal:2006tn,Balasubramanian:2004uy} (inherently liftable to dS by the pure moduli sector itself) discussed here. These three classes of moduli stabilizing vacua are three branches of solutions in the same low-energy 4d ${\\cal N}=1$ supergravity arising from type IIB compactified on a Calabi-Yau orientifold with D7-branes.  In section~\\ref{T_stab}, we will review the method of `K\\\"ahler uplifting' and analytically derive the existence of the meta-stable dS vacuum for the volume modulus of a one-parameter Calabi-Yau compactification with $h^{1,1}=1$ K\\\"ahler modulus, and then extend this to the case of several K\\\"ahler moduli $h^{1,1}>1$ explicitly. The interplay of perturbative and non-perturbative effects implies for $h^{1,1}=1$ that here a structure of \\emph{two} terms with alternating signs is sufficient to approximate the volume modulus scalar potential and its tunable dS vacuum. This contrasts with the `3-term structure' generically necessary in purely perturbatively stabilized situations~\\cite{Saltman:2004jh,Silverstein:2007ac}. For $h^{1,1}>1$ a `3-term structure' reappears for the additional $h^{1,1}-1$ blow-up K\\\"ahler moduli of a `swiss cheese' Calabi-Yau.  Finally, we will show that we can express the existence of the meta-stable dS vacuum for the volume modulus in terms of a \\emph{sufficient} condition on the microscopic parameters. These are consisting of the fluxes, the D7-brane configuration, and the Euler number of the Calabi-Yau governing the perturbative ${\\cal O}(\\alpha'^3)$-correction, which are all in turn determined by the underlying F-theory compactification on an elliptically fibred Calabi-Yau fourfold. Thus, the result amounts to a sufficient condition for the existence of meta-stable dS vacua in terms of purely F-theory geometric and topological data which can be satisfied for a sizable subclass of all 4d ${\\cal N}=1$ F-theory compactifications, instead of just single `lamp post' models. We also check that our sufficient condition satisfies the necessary condition for meta-stable dS vacua in 4d ${\\cal N}=1$ supergravity given in~\\cite{Covi:2008ea} and the longevity of the metastable vacuum under tunneling.  Section~\\ref{TS_stab} includes the dilaton into a full analytical treatment of the combined dS minimum. We show that supersymmetry breaking happens predominantly in the volume modulus direction, and explicitly determine the shift of the dilaton away from its flux-stabilized supersymmetric locus as suppressed by inverse powers of the volume of the Calabi-Yau.  Section~\\ref{TSU_stab_gen} extends the analysis by including an arbitrary number of complex structure moduli with unspecified dependence in the K\\\"ahler and superpotential. We then show that the shift of the complex structure moduli and the dilaton in general is suppressed by inverse powers of the volume, and that the dilaton and all complex structure moduli generically are fixed at positive-definite masses. Finally, we estimate the backreaction of the shifted dilaton and complex structure moduli onto the volume modulus. The ensuing shift of the stabilized volume is generically found to be small and suppressed by inverse powers of the volume. This crucially extends the sufficient condition for the existence of dS vacua in type IIB F-theory compactifications to a large class of `swiss cheese' style fluxed Calabi-Yau compactifications with arbitrary $h^{1,1}<h^{2,1}$.  In section~\\ref{TSU_stab}, we apply our methods to a simple toy model where the K\\\"ahler and superpotential of complex structure moduli are approximated by the structure found in a torus compactification. We verify the general results of the previous sections, and show that the shifts of the moduli and the backreaction effects are either independent of the number of complex structure moduli $h^{2,1}$, or decreasing as an inverse power of $h^{2,1}$. We conclude in section~\\ref{Concl}.  While this paper was being finished, we became aware of~\\cite{deAlwis:2011dp}, whose section 2 contains overlapping results with our section~\\ref{T_stab}. The main results of section~\\ref{T_stab} and~\\ref{TS_stab} here have first been presented in talk by one of the authors in~\\cite{SFBmeeting}. Additionally, we find numerical disagreement concerning the values of $x$ in section~\\ref{T_stab} permissible for a meta-stable dS vacuum of $T$ compared to the results for the same quantity given in section 2 of~\\cite{deAlwis:2011dp} due to an approximation used between eq.s (16) and (17) \\emph{ibid}. ",
        "conclusions": "`K\\\"ahler uplifting' has the benefit of generating meta-stable dS vacua in terms of just background 3-form fluxes, D7-branes and the leading perturbative ${\\cal O}(\\alpha'^3)$-correction, which data are completely encoded in terms of the underlying F-theory compactification on a fluxed Calabi-Yau fourfold. In addition, supersymmetry is spontaneously broken (typically $\\sim M_{\\rm GUT}$ here, which is below the KK scale) by an F-term generated in the volume moduli sector. No extra anti-branes, D-terms, or F-term generating matter fields are needed or involved. We emphasize that the perturbative ${\\cal O}(\\alpha'^3)$-correction breaks supersymmetry spontaneously. In a similar fashion as in spontaneously broken gauge theories where the ground state does not respect the full symmetry but the underlying action is invariant and thus protected from dangerous corrections by the full gauge symmetry, we expect the spontaneous breaking of supersymmetry to enhance control over the back-reaction of the various branes and orientifolds. This is because the underlying action itself still does respect the full unbroken supersymmetry. Furthermore, the backreaction of the supersymmetric D7-branes with their gauge groups is fully accounted for in the language of F-theory describing them as ADE-type singularities of the elliptic fibration of the F-theory 4-fold.  Here, we have developed a method towards a rigorous analytical understanding of `K\\\"ahler uplifting' driven by the leading ${\\cal O}(\\alpha'^3)$ correction to the K\\\"ahler potential of the volume moduli. Our derivation was carried out in the presence of an arbitrary number $h^{2,1}$ of complex structure moduli. A large value of 3-cycles $h^{2,1}={\\cal O}(100)$  is a prerequisite to use the associated 3-form fluxes for the required fine-tuning of the cosmological constant. We have given an argument for the existence of meta-stable dS vacua in so-called `swiss cheese' type Calabi-Yau compactifications of negative Euler characteristic with an arbitrary number $h^{1,1}$ of K\\\"ahler moduli. The interplay of perturbative and non-perturbative effects in generating this dS minimum implies for one-parameter models with $h^{1,1}=1$ that here a structure of \\emph{two} terms with alternating signs is sufficient to approximate the volume modulus scalar potential and its tunable dS vacuum. This contrasts with the `3-term structure' generically necessary in purely perturbatively stabilized situations~\\cite{Saltman:2004jh,Silverstein:2007ac}. For $h^{1,1}>1$ a `3-term structure' reappears for the additional $h^{1,1}-1$ blow-up K\\\"ahler moduli of a `swiss cheese' Calabi-Yau.   Exploiting this `2-term structure' (or, alternatively, the `3-term structure' for $h^{1,1}>1$), we have shown that we can express the existence of the meta-stable dS vacuum for the volume modulus in terms of a \\emph{sufficient} condition on the microscopic parameters. These are consisting of the fluxes, the D7-brane configuration, $h^{1,1}$, and the Euler number of the Calabi-Yau governing the perturbative ${\\cal O}(\\alpha'^3)$-correction, which are all in turn determined by the underlying F-theory compactification on an elliptically fibred Calabi-Yau fourfold. Thus, the result amounts to a sufficient condition for the existence of meta-stable dS vacua in terms of purely F-theory geometric and topological data which can be satisfied for a sizable subclass of all 4d ${\\cal N}=1$ F-theory compactifications, instead of just single `lamp post' models.  Our sufficient condition survives both explicit inclusion of dilaton stabilization by fluxes as well as an arbitrary number of complex structure moduli. Supersymmetry breaking happens predominantly in the volume modulus direction, and explicitly determine the shift of the dilaton and all complex structure moduli away from their flux-stabilized supersymmetric locus as suppressed by inverse powers of the volume of the Calabi-Yau. We have also checked the longevity of the metastable vacuum under tunneling.  However, there are still some possible caveats. Possible mixing with KK modes may not decouple fast enough with increasing volume of the Calabi-Yau in a given example to avoid further tachyonic directions. The structure of the 1-loop determinants of the condensing gauge groups used for volume moduli stabilization is not known in general, and it may display a dependence on the complex structure moduli which might be sufficiently strong to possibly derail our perturbative treatment of complex structure stabilization in some examples.  Finally, we have estimated the backreaction of the shifted dilaton and complex structure onto the volume modulus. The ensuing shift of the stabilized volume is generically found to be small and suppressed by inverse powers of the volume.  In conclusion, we have given arguments towards a \\emph{sufficient} condition for the existence of meta-stable dS vacua in terms of, ultimately, purely F-theory data which can be satisfied for the sizable class of fluxed `swiss cheese' type Calabi-Yau compactifications with arbitrary $h^{1,1}<h^{2,1}$."
    },
    "1107/1107.2924_arXiv.txt": {
        "abstract": "We report ground-based follow-up observations of the exceptional source, ID\\,141, one the brightest sources detected so far in the {\\it H}-ATLAS cosmological survey. ID\\,141 was observed using the IRAM 30-meter telescope and Plateau de Bure interferometer (PdBI), the Submillimeter Array (SMA) and the Atacama Pathfinder Experiment (APEX) submillimeter telescope to measure the dust continuum and emission lines of the main isotope of carbon monoxide and carbon ([C\\,{\\small I}] and [C\\,{\\small II}]). The detection of strong CO emission lines with the PdBI confirms that ID\\,141 is at high redshift ($z=4.243 \\pm 0.001$).  The strength of the continuum and emission lines suggests that ID\\,141 is gravitationally lensed.  The width ($\\Delta V_{\\rm FWHM} \\sim 800 \\, \\rm km \\, s^{-1}$) and asymmetric profiles of the CO and carbon lines indicate orbital motion in a disc or a merger. The properties derived for ID\\,141 are compatible with a ultraluminous ($L_{\\rm FIR} \\sim 8.5\\pm0.3 \\times 10^{13} \\, \\mu_{\\rm L}^{-1} \\, \\rm L_\\odot$, where $\\mu_L$ is the amplification factor), dense ($n \\approx 10^4 \\, \\rm cm^{-3}$) and warm ($T_{\\rm kin} \\approx 40 \\, \\rm K$) starburst galaxy, with an estimated star-formation rate of (0.7 to 1.7)$\\times 10^4 \\, \\mu_{\\rm L}^{-1} \\rm \\, M_\\odot / yr$. The carbon emission lines indicate a dense ($n \\approx 10^4 \\rm cm^{-3}$) Photo-Dominated Region, illuminated by a far-UV radiation field a few thousand times more intense than that in our Galaxy. In conclusion, the physical properties of the high-$z$ galaxy ID\\,141 are remarkably similar to those of local ultraluminous infrared galaxies. ",
        "introduction": "Over the last decade, submillimeter surveys have revolutionized our understanding of the formation and evolution of galaxies by uncovering a population of high-redshift, dust-obscured systems that are forming stars at a tremendous rate (e.g.\\ Smail et al.\\ 1997; Hughes et al.\\ 1998; Blain et al.\\ 2002).  Submillimeter galaxies (SMG) are high-redshift analogues of the local ultraluminous infrared galaxies (ULIRGs) in terms of their CO, radio and infrared properties (Tacconi et al.\\ 2008, 2010). The general importance of ultraluminous sources for galaxy evolution has been highlighted by recent {\\it Spitzer} and {\\it Herschel} results that show that these sources contribute significantly to the total amount of star formation in the early universe \\cite{Magnelli2009}.  Until recently, observations at submillimeter wavelengths have been conducted mainly from the ground and/or from balloon experiments and covered only relatively small areas of the sky. The discovery rate of rare objects such as strongly lensed sources was therefore limited. Nevertheless, due to selection effects, many of the known high-redshift submillimeter sources are gravitationally lensed, such as {\\it IRAS}\\,F10214+4724, which is amplified by a factor of $\\sim 10\\times$ \\cite{Rowan-Robinson1991, Downes1995}.  However, these limited surveys found few strongly lensed sources.  This situation changed dramatically with the 87\\,deg$^2$ deep field survey at wavelengths of 1.4 and 2.0~mm taken with the South Pole Telescope \\cite{Vieira2010} and, since then, the {\\it Herschel} space observatory \\cite{Pilbratt2010} has been able to map an enormously extended area of the sky at submillimeter wavelengths.  In particular, the wide-area {\\it Herschel} Astrophysical Terahertz Large Area Survey or {\\it H}-ATLAS \\cite{Eales2010} with the SPIRE \\cite{Griffin2010} and the PACS \\cite{Poglitsch2010} instruments will ultimately map a total of about 570 square-degrees to around (or below) the confusion limit in five bands from 100 to 500~$\\rm \\mu m$. In addition to uncovering hundreds of thousands of SMGs, the {\\it Herschel} surveys will enable discovery of the brightest infrared sources in the universe, including many strongly lensed sources because of the high magnification bias in the submillimeter bands \\cite{Blain1996}. These surveys and samples of lensed sources are unique and will remain so for the foreseeable future.  Negrello et al.\\ (2007) predicted that most of the brightest 500-$\\rm \\mu m$ sources (i.e.\\ those with $S_{\\rm 500 \\, \\mu m} > 100$\\,mJy) detected at high redshift by the {\\it Herschel} space observatory will consist of strongly lensed submillimeter galaxies, a population of local star-forming galaxies -- easily identified in shallow Sloan Digital Sky Survey (SDSS) images -- a relatively small fraction of radio-bright AGNs -- easily identifiable in shallow radio surveys \\cite{Negrello2010} -- as well as nearby spirals.  The ongoing surveys have confirmed this prediction with the detection of submillimeter-bright ($S_{\\rm 500 \\, \\mu m} > 100$\\,mJy) sources that have been shown, based on ground-based follow-up observations, to be gravitationally lensed SMGs in the distant ($z \\sim 1.5-4$) Universe \\cite{Negrello2010, Conley2011}.  The lensing boosts the sensitivity of observations and improves their spatial resolution enabling detailed studies of the populations responsible for the bulk of the background \\cite{Blain1999, Hopwood2010}.  Follow-up observations with ground-based telescopes, especially by (sub)millimeter facilities are essential to further probe the physical characteristics of these outstanding sources.  A recent and striking example of a lensed source discovered in the pre-{\\it Herschel} era in the submillimeter is the $z=2.3$ galaxy SMM\\,J2135$-$0102 \\cite{Swinbank2010} or the `Cosmic Eyelash'. This typical SMG, which is amplified by a factor of $32.5\\times$, has been observed in exquisite detail enabling diagnosis of the properties of its molecular and atomic gas \\cite{Danielson2011, Swinbank2011}.  Here, we report ground-based follow-up observations of an exceptional source uncovered by the {\\it H}-ATLAS survey, HATLAS\\,J142413.9+022304 (hereafter ID\\,141), including a measurement of its redshift and a derivation of the properties of its molecular and atomic gas content.  Throughout the paper we use a $\\Lambda$CDM cosmology with $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m=0.27$ and $\\Omega_{\\Lambda}=1-\\Omega_m$ \\cite{Spergel03}. ",
        "conclusions": "We have used four ground-based (sub)millimeter facilities -- the IRAM 30-meter and PdBI, the SMA and APEX -- to follow up the exceptional source ID\\,141 that is the strongest 500-$\\rm \\mu m$ riser yet discovered in the {\\it Herschel} cosmological surveys. The detection of three bright CO emission lines using the PdBI confirms that the source is at high redshift ($z=4.243$). The strength of the continuum and the intensities of the molecular and atomic emission lines of ID\\,141 suggest the source is gravitationally lensed, in line with what was already found for other sub-millimeter bright, high-redshift, galaxies detected with the {\\it Herschel} space observatory \\cite{Negrello2007, Negrello2010}. The magnification factor is still unknown but is likely to be in the range $10 < \\mu_{\\rm L} <30$ as found for other gravitationally lensed SMGs uncovered in the {\\it H}-ATLAS survey.  The properties of the molecular and atomic gas in ID\\,141 indicate properties of a luminous ($L_{\\rm FIR} \\sim 3 - 8 \\times 10^{12} \\, L_\\odot$), dense ($n \\approx 10^4 \\, \\rm cm^{-3}$) and warm ($T_{\\rm kin} \\approx 40 \\rm \\, K$) starburst galaxy that are comparable to those derived for local starburst galaxies or high-redshift submillimeter galaxies. The asymmetric double profiles of the carbon and CO emission lines are indicative of orbital motions in a disc or a merger. The source is barely resolved in the present data indicating a physical size that is smaller than $14/\\mu_{\\rm L}$~kpc, corresponding to an upper limit of $2^{\\prime\\prime}$.  The measurements of the carbon lines ([C\\,{\\small I}] and [C\\,{\\small II}]) in ID\\,141 enable a first study of the gas associated with the PDR in ID\\,141 and analysis of the relative cooling of the carbon and CO emission in this galaxy, that is dominated by that of $\\rm C^+$. The [C\\,{\\small II}]-to-FIR luminosity ratio, $(7.3 \\pm 1.3) \\times 10^{-4}$, is larger than the one found for other galaxies of comparable luminosity and redshift, and consistent with recent findings that the scatter in the $L_{\\rm C{\\small II}}/L_{\\rm FIR}$ versus $L_{\\rm FIR}$ diagram is indeed large. With a ratio $L_{\\rm FIR}/M_{\\rm H_2} \\sim 230 \\, \\rm L_\\odot \\, M_\\odot^{-1}$, ID\\,141 follows the trend reported by Graci\\'a-Carpio et al.\\ (2011) in the $L_{\\rm C{\\small II}}/L_{\\rm FIR}$ versus $L_{\\rm FIR}/M_{\\rm M_{H_2}}$ plane, suggesting both a high ionization parameter and the fact that ID\\,141 could be a merger.  The present observations underline the importance of observations done at frequencies in between 85 and 360~GHz to study the properties of galaxies at redshift $z>4$. In particular, the lower frequencies (in the 3 and 2~mm windows) are essential to probe the peak of the CO emission that, for such starburst galaxies, lies around $J_{\\rm upper}$ 4 to 5, as well as the two emission lines of [C\\,{\\small I}] or molecules other than CO such as water. Follow-up observations of very-high-redshift, obscured galaxies found in future deep surveys, especially the less luminous objects, will need sensitive interferometers operating at frequencies below 350~GHz such as ALMA, the EVLA or the future upgrade of the Plateau de Bure interferometer (NOEMA) in order to fully explore the properties of the molecular and atomic gas in starburst galaxies that were active when the universe was about 1~Gyr old."
    },
    "1107/1107.5390_arXiv.txt": {
        "abstract": "We consider the modified restricted three body problem with power-law density profile of disk, which rotates around the center of mass of the system with perturbed mean motion. Using analytical and numerical methods we have found equilibrium points and examined their linear stability. We have also found the zero velocity surfaces for the present model. In addition to five equilibrium points there is a new equilibrium point on the line joining the two primaries. It is found that $L_2$ and $L_3$ are stable for some values of inner and outer radius of the disk while collinear points are unstable, but $L_4$ is conditionally stable for mass ratio less than that of Routh's critical value. Lastly we have obtained the effects of radiation pressure, oblateness and mass of the disk. ",
        "introduction": "The problem, after imposing a restriction as one body of the three body problem is of an infinitesimal(negligible) mass and remaining  other two are of finite masses, is known as restricted three body problem(RTBP). The governing force of motion of the RTBP is mainly the gravitational forces exerted by the finite masses also known as primaries. In the RTBP if we take bigger primary as a source of radiation then problem called as photogravitational RTBP which is generalized by taking smaller primary as an oblate spheroid.  The Chermnykh-like problem which was first time studied by \\cite{Chermnykh1987VeLen.......73C}, deals the motion of an infinitesimal mass in the orbital plane of a disk which rotates around the center of mass of the primaries with constant angular velocity $n$. \\cite{K.Gozdzieski1998CeMDA..70...41G}, examined the problem in the sense of nonlinear stability of equilibrium points and also obtained the range of parameter for the same.   The Chermnykh-Like problem has a number of applications in different areas as celestial mechanics (\\cite{Gozdziewski1999CeMDA..75..251G}), chemistry (\\cite{Strand1979JChPh..70.3812S}) etc. Also, the importance of the problem have seen in the extra solar planetary system (\\cite{Rivera2000ApJ...530..454R}, \\cite{Jiang2001AA...367..943J} etc.).  Further the effect of disc is very helpful in the study of resonance capture of Kuiper Belt Objects (KBOs) as given in \\cite{Jiang2004MNRAS.355L..29J}. \\cite{Papadakis2005Ap&SS.299..129P}, by taking assumptions as constant mass parameter and variable angular velocity parameter, analyzed the equilibrium point and zero velocity curve; \\cite{Papadakis2005Ap&SS.299...67P} also studied the problem numerically. \\cite{Jiang2006Ap&SS.305..341J} examined the Chermnykh problem with {$\\mu$} = 0.5 and shown a deviation in the result of classical RTBP; also they have found the new equilibrium points in spite of Lagrangian points.  \\cite{YehLCJiang2006} have found the condition of existence of new equilibrium points analytically and numerically.   \\cite{IshwarKushvah2006math......2467I} examined the linear stability of triangular points with P-R drag. Again \\cite{Kushvah2008Ap&SS.318...41K} examined the stability of collinear points and found  unstable points.    Motivating by the importance and applications of the Chermnykh-like problem, we modeled a generalized photogravitational Chermnykh-like problem (section-\\ref{sec:model}) in which we consider the angular velocity parameter $n>1$ which depends on the gravitational force of the disk assumed, radiation force of radiating (bigger) primary and oblateness factor of smaller primary which is an oblate spheroid.Further we determine the equilibrium points (section-\\ref{sec:equipts}) and zero velocity surface (section-\\ref{sec:zvs}) and then find the stability of the equilibrium points (section-\\ref{sec:stb}) finally conclude the results (section-\\ref{sec:conclude}). ",
        "conclusions": "\\label{sec:conclude} We have studied the dynamical properties of modified restricted three body problem and found that there exists a new equilibrium point in addition to five equilibrium points in classical problem. The zero velocity surface have obtained and examined the stability of equilibrium points. It is found that the $L_2$ and $L_3$ are stable for certain values of inner and outer radius of the disk and other collinear points are unstable while $L_4$ is conditionally stable upto the Routh's value of mass ratio. Thus stability region spans with the width of the disk. Hence, we conclude that the shape and size of the disk are  very significant for the motion of the bodies in space."
    },
    "1107/1107.3749_arXiv.txt": {
        "abstract": "We investigate here the genericity and stability aspects for naked singularities and black holes that arise as the final states for a complete gravitational collapse of a spherical massive matter cloud. The form of the matter considered is a general {\\it Type I} matter field, which includes most of the physically reasonable matter fields such as dust, perfect fluids and such other physically interesting forms of matter widely used in gravitation theory. We first study here in some detail the effects of small pressure perturbations in an otherwise pressure-free collapse scenario, and examine how a collapse evolution that was going to the black hole endstate would be modified and go to a naked singularity, once small pressures are introduced in the initial data. This allows us to understand the distribution of black holes and naked singularities in the initial data space. Collapse is examined in terms of the evolutions allowed by Einstein equations, under suitable physical conditions and as evolving from a regular initial data. We then show that both black holes and naked singularities are generic outcomes of a complete collapse, when genericity is defined in a suitable sense in an appropriate space. ",
        "introduction": "The aim of the present work is twofold. Firstly, we study here the effect of introducing small pressure perturbations in an otherwise pressure-free gravitational collapse which was to terminate in a black hole final state. For such a purpose, spherically symmetric models of black hole and naked singularity formation for a general matter field are considered, which undergo a complete gravitational collapse under reasonable physical conditions while satisfying suitable energy conditions. Secondly, we investigate the genericity and stability aspects of the occurrence of naked singularities and black holes as collapse endstates. The analysis of pressure perturbations in known collapse models, inhomogeneous but otherwise pressure-free, shows how collapse final states in terms of black hole or naked singularity are affected and altered. This allows us to examine in general how generic these outcomes are and we study in the initial data space the set of conditions that lead the collapse to a naked singularity and investigate how `abundant' these are. While it is known now for some time that both black holes and naked singularities do arise as collapse endstates under reasonable physical conditions, this helps us understand and analyze in a clear manner the genericity aspects of occurrence of these objects in a complete gravitational collapse of a massive matter cloud in general relativity.    The physics that is accepted today as the backbone of the general mechanism describing the formation of black holes as the endstate of collapse relies on the very simple and widely studied Oppenheimer-Snyder-Datt (OSD) dust model, which describes the collapse of a spherical cloud of homogeneous dust \\cite{OSD}, \\cite{Datt}. In the OSD case, all matter falls into the singularity at the same comoving time while an horizon forms earlier than the singularity, thus covering it. A black hole results as the endstate of collapse. Still, homogeneous dust is a highly idealized and unphysical model of matter. Taking into account inhomogeneities in the initial density profile it is possible to show that the behaviour of the horizon can change drastically, thus leaving two different outcomes as the possible result of generic dust collapse: the black hole, in which the horizon forms at a time anteceding the singularity, and the naked singularity, in which the horizon is delayed thus allowing null geodesics to escape the central singularity where the density and curvatures diverge, to reach faraway observers \\cite{dust1}-\\cite{dust5}.    It is known now that naked singularities do arise as a general feature in General Relativity under a wide variety of circumstances. Many examples of singular spacetimes can be found, but their relevance in models describing physically viable scenarios has been a matter of much debate since the first formulation of the Cosmic Censorship Hypothesis (CCH) \\cite{Penrose}. In particular, the formation of naked singularities in dynamical collapse solutions of Einstein field equations remains a much discussed problem of contemporary relativity. The CCH, which states that any singularity occurring in the universe must be hidden within an event horizon and therefore not visible to faraway observers, has remained at the stage of a conjecture for more than four decades now. This is also because of the difficulties lying in a concrete and definitive formulation of the conjecture itself. While no proof or any mathematically rigorous formulation of the same exists in the context of dynamical gravitational collapse (while some proofs exist for particular classes of spacetimes that do not describe gravitational collapse, as in \\cite{Dafermos1} and \\cite{Dafermos2}), many counterexamples have been found over the past couple of decades \\cite{Ref1}-\\cite{Ref6}. Many of these collapse scenarios are restricted by some simplifying assumptions such as the absence of pressures (dust models) or the presence of only tangential pressures \\cite{tang1}-\\cite{tang7}. It is well-known that the pressures cannot be neglected in realistic models describing stars in equilibrium. It seems natural therefore that if one wishes to study analytically what happens during the last stages of the life of a massive star, when its core collapses under its own gravity thus forming a compact object as a remnant, pressures must be taken into account.    Therefore, further to early works that showed the occurrence of naked singularities in dust collapse, much effort has been devoted to understanding the role played by pressures \\cite{press1}-\\cite{press7}. The presence of pressures is a crucial element towards the description of realistic sources as we know that stars and compact objects are generally sustained by matter with strong stresses (either isotropic or anisotropic). At first it was believed that the naked singularity scenario could be removed by the introduction of pressures, thus implying that more realistic matter models would lead only to the formation of a black hole. We now know that this is not the case. The final outcome of collapse with pressure is entirely decided by its initial configuration and allowed dynamical evolutions and it can be either a black hole or a naked singularity. Furthermore it is now clear that within spherical collapse models (be it dust, tangential pressure or others) the data set leading to naked singularities is not a subset of `zero measure' of the set of all possible initial data.   Despite all this work we can still say that much more is to be understood about the role that general pressures play during the final stages of collapse. Perfect fluid collapse has been studied mostly under some simplifying assumptions and restrictions in order to gain an understanding, but a general formalism for perfect fluids described by a physically valid equation of state is still lacking due to the intrinsic difficulties arising from Einstein equations. Considering both radial and tangential pressures is a fundamental step in order to better understand what happens in the ultra-dense regions that forms at the center of the collapsing cloud prior to the formation of the singularity. For this reason, perfect fluids appear as a natural choice since these are the models that are commonly used to describe gravitating stars in equilibrium and since it can be shown that near the center of the cloud regularity implies that matter must behave like a perfect fluid.   In the present paper we use a general formalism to analyze the structure of collapse in the presence of perfect fluid pressures. This helps to understand better realistic collapse scenarios and their outcomes and brings out clearly the role played by pressures towards the formation of black holes or naked singularities as the endstate of collapse. We examine what are the key features that determine the final outcome of collapse in terms of a black hole or a naked singularity when perfect fluids, without any restriction imposed by the choice of an equation of state, are considered. The reason we do not assume an explicit equation of state here is that the behavior of matter in ultra-dense states in the final stages of collapse is unknown. On the other hand, having regularity and energy conditions satisfied provides a physically reasonable framework to study collapse endstates.   We find that not only naked singularities are not ruled out in perfect fluid collapse scenarios but also that the separation between the black hole region and the naked singularity region in the space of all possible evolutions has some interesting features. In particular, we show that the introduction of small pressures can drastically change the final fate of the well-known pressureless models. For example, we see that adding a small pressure perturbation to an inhomogeneous dust model leading to a black hole can be enough to change the outcome of collapse to a naked singularity, and viceversa.   Further to this, we investigate here the space of initial data and collapse evolutions in generality, in order to examine the genericity of naked singularities in collapse. To study small pressure perturbations as well as the genericity and stability aspects, we use the general formalism for spherically symmetric collapse developed earlier \\cite{ndim1}, \\cite{ndim2} in order to address the basic problem of how generic is a given collapse scenario which leads to the formation of naked singularities. Given the existence of an increasing number of models describing collapse leading to a naked singularity, the issue of genericity and stability of such models in the space of initial data has become the crucial ingredient in order to decide whether the Cosmic Censorship Hypothesis in its present form should be conserved, modified or dropped altogether.    It should be noted, however, that the concepts such as genericity and stability are far from well-defined in a unique manner in general relativity, as opposed to the Newtonian gravity. A major difficulty towards such a task is the non-uniqueness of topology, or the concept of `nearness' itself in a given spacetime geometry \\cite{HE}. One could define topology on a space of spacetime metrics by requiring that the metric component values are `nearby' or also additionally requiring that their $n$-th derivatives are also nearby, and in each case the resulting topologies will be different. This is in fact connected in a way with the basic problem in arriving at a well-formulated statement of the cosmic censorship itself. There have been attempts in the past to examine the genericity and stability of naked singularities in special cases. For example, in \\cite{C} it was shown that for certain classes of massless scalar field collapse the initial data leading to naked singularity has, in a certain sense, a positive codimension, and so the occurrence of naked singularity is unstable in that sense. On the other hand, it was shown in \\cite{Sar-Pramana}, \\cite{SG} that naked singularity occurrence is stable in the sense of the data sets leading to the same being open in the space of initial data. But these need not be dense in this space and so `non-generic' if we use the definition of `genericity' in the sense given in the dynamical systems theory (where a set of initial data leading to a certain outcome is said to be generic if it is open and dense within the set of all initial data). In that case, however, both black hole and naked singularity final states turn out to be `non-generic'.   Therefore, in the following we adopt a more physical definition of `genericity' in the sense of `abundance', and we call generic an initial data set that has a non-zero measure, and which is open (though not necessarily dense) in the set of all initial data. With this definition, the results obtained in \\cite{Sar-Pramana} and \\cite{SG} would mean that both black hole and naked singularity are generic collapse endstates. We note that we do not deal here with the self-similar models, or scalar fields, which is a somewhat special case. Therefore, the issue of genericity and stability of naked singularities in collapse remains wide open, for spherically symmetric as well as non-spherical models and for different forms of matter fields. Our consideration here treats in this respect a wide variety of physically reasonable matter fields for spherically symmetric gravitational collapse.    In section \\ref{einstein}, the general structure for Einstein equations to study spherical collapse is reviewed and we describe how the equations can be integrated thus obtaining the equation of motion for the system. In Section \\ref{perturbation} we examine the structure of the initial data sets of collapse leading to black hole and naked singularity to gain an insight on genericity of such outcomes for some special models and effect of introducing small pressure perturbations is investigated. Section \\ref{genericity} then considers the genericity aspects of the outcomes of collapse with respect to initial data sets. We prove that the initial data sets leading to black holes and naked singularities in the space of all initial data sets for perfect fluid collapse are both generic. Section \\ref{eos} is devoted to a brief discussion on equations of state. Finally, in section \\ref{remarks} we outline the key features of the above approach and its advantages, and point to possible future uses of the same for astrophysical and numerical applications. ",
        "conclusions": "\\label{remarks}  We have studied here the general structure of complete gravitational collapse of a sphere composed of perfect fluid without a priori requiring an equation of state for the matter constituents, thus allowing for the freedom to choose the mass function arbitrarily, as long as physical reasonableness as imposed by regularity and energy conditions is satisfied.   The interest of such an analysis comes from the fact that the class of perfect fluid models for matter is considered to be physically viable for the description of realistic objects in nature such as massive stars and their gravitational collapse. Typically perfect fluids are considered to be physically more sound than models where matter is approximated by dust-like behaviour {\\it i.e.} without pressures, or where matter is sustained by only tangential pressures (though the `Einstein cluster' describing a spherical cloud of counter-rotating particles has been shown to have some non-trivial  physical validity \\cite{cluster1}-\\cite{cluster3}). What we have shown is that, within the class of perfect fluid collapses, both final outcomes, namely black holes and naked singularities, can be equally possible depending on the choice of the initial data and the free function $F$. In fact our results show that naked singularities and black holes are both possible final states of collapse, much in the same way as it has already been proven in the simpler cases of inhomogeneous dust and matter exhibiting only tangential stresses. The sets of initial data leading to either of the outcomes share the same properties in terms of genericity and stability.     The structure of initial data sets in the case of OSD, LTB and pressure collapse and their inter-relationship is, however, a complicated issue. Nevertheless, we can comment on this based on the studies in \\cite{Sar-Pramana}, \\cite{SG}, and the results proved in sections \\ref{perturbation} and \\ref{genericity} in this paper. In \\cite{SG} it was proven that the space of initial data $\\{ M(r), b( r )\\}$ leading LTB collapse to black hole or naked singularity forms an open subset of $X \\times X$, where $X$ is the infinite dimensional Banach space of real $\\mathcal{C}^1$ functions defined on the domain. As per the analogous results in the case of non-vanishing tangential pressures it follows that the initial data set leading the collapse to OSD black holes is a non-generic subset of space of $\\{  M(r), b( r )\\}$. The shortcoming of the tangential pressure case being that is not wholly physically satisfactory. Therefore to investigate the perfect fluid case, we moved to a `bigger' space $X \\times X \\times X$, since the initial data set comprises of $\\{M(r, v_i), p (r, v_i), b( r)\\}$. Thus, in this space, the initial data set $I_{OSD}$ or $I_{TBL}$ or the union of both the sets, will become non-generic. Mathematically speaking, this set is meager or nowhere dense in the space $X \\times X \\times X$. This is proved by the study performed here. Thus, the initial configurations for the end states in the case of LTB or OSD models lie on the critical surface separating the two possible outcomes of collapse as discussed above.    This analysis in fact shows how the structure of Einstein equations is very rich and complex, and how the introduction of pressures in the collapsing cloud opens up a lot of new possibilities that, while showing many interesting dynamical behaviours, do not rule out either of the two possible final outcomes.      There are physical reasons, however, for the perfect fluid model to be subdued to the choice of an equation of state and there is also increasing evidence that such an equation of state cannot hold during the whole duration of the dynamical collapse. In fact there are indications that as the collapsing matter approaches the singularity large negative pressures arise, thus making the equation of state relating density to pressures depart from usual well-known equations of stellar equilibrium. Nevertheless the study of similar scenarios with linear or polytropic equations of state can give insights in the initial stages of collapse of a star. All this is very important from astrophysical point of view where still little is known of the processes that happen towards the very end of the life of a star, when in a catastrophic supernova explosion the outer layers are expelled and the inner core collapses under its own gravity.    As we mentioned, due to the intrinsic complexity of Einstein equations for perfect fluid collapse, it is generally possible to solve the system of equations only under some simplifying assumptions (like the choice of a specific mass function), and only close to the center of the cloud. The indications provided by the present analysis are then a first step towards a better understanding of what happens in the last stages of the complete gravitational collapse of a realistic massive body.    Furthermore the above formalism could possibly be used as the framework upon which to develop possible numerical simulations of gravitational collapse. As seen in the comoving frame, the positivity of $\\chi_1$ (or $\\chi_2$)) is the necessary and sufficient condition for the singularity to be visible, at least locally. Numerical models of a collapsing star made of a perfect fluid with a polytropic equation of state (or a varying equation of state that takes into account the phase transitions that occur in matter under strong gravitational fields), with the addition of rotation and possibly electromagnetic field might help us better understand whether the inner ultradense region that forms at the center of the collapsing cloud when the apparent horizon is delayed, might be visible globally and have some effects on the outside universe. Many numerical models that describe dynamical evolutions leading to the formation of black holes exist both in gravitational collapse as in merger of compact objects such as neutron stars (see e.g. \\cite{Rezzolla1}, \\cite{Rezzolla2}). but a fully comprehensive picture of what happens in the final moments of the life of a star is still far away.    Close to the formation of the singularity, gravitationally repulsive effects, possibly due to some quantum gravitational corrections, are likely to take place. If such phenomena can interact with the outer layers of the collapsing cloud they might create a window to the physics of high gravitational fields whose effects might be visible to faraway observers. This scenario might in turn imply the visibility of the Planck scale physics or new physics close to the singularity, the presence of a quantum wall that might cause shock-waves from within the Schwarzschild radius that might give rise to different type of emissions and explosions, with photons or high energy particles escaping from the ultradense region. Collisions of particles with arbitrarily high center of mass energy near the Cauchy horizon might also happen \\cite{accel}.     Overall, the analysis led over the past few years seems to suggest that the Oppenheimer-Snyder-Datt scenario is indeed too restrictive to account for the richness of realistic dynamical models in general relativity. The occurrence of naked singularities in gravitational collapse appears to be a well established fact and a lot of intriguing new physics might arise from the future study of more detailed collapse models."
    },
    "1107/1107.3113_arXiv.txt": {
        "abstract": "We present high spatial and temporal resolution \\ca\\ observations of a kink wave in an on-disk chromospheric active region fibril. The properties of the wave are similar to those observed in off-limb spicules. From the observed phase and period of the wave we determine a lower limit for the field strength in the chromospheric active region fibril located at the edge of a sunspot to be a few hundred Gauss. We find indications that the event was triggered by a small-scale reconnection event higher up in the atmosphere. ",
        "introduction": "Recent observations have strengthened the connection between spicules and coronal heating, either via spicules supplying hot plasma to the corona \\citep{DePontieu+others2011} or energy dissipation from MHD waves propagating along spicules (e.g.,~\\citealt{Kukhianidze+others2006,DePontieu+others2007,He+others2009}). The first observations of transverse oscillations, indicative of MHD waves, in spicules were made already in the late 1960s \\citep{Pasachoff+others1968}. Since the discovery that MHD waves are fairly ubiquitous in spicules \\citep{DePontieu+others2007} the interest has been renewed. For a recent observational review of spicule oscillations see \\cite{Zaqarashvili+Erdelyi2009}. Several authors, e.g., \\cite{DePontieu+others2007,He+others2009}, have identified the waves as kink or Alfv\\'en and estimated that they carry enough energy to heat the corona.  Based on theoretical considerations we expect to find four different modes of waves in cylindrical plasma geometries appropriate for solar conditions: kink, sausage, longitudinal, and torsional modes. The first three modes are compressible, while the fourth one is not. All of the modes have been observed in the solar corona (although the identification of the pure, and thus incompressible, Alfv\\'en waves has been questioned by \\citealt{VanDoorsselaere+others2008}). Alfv\\'en and kink waves can be triggered by granular buffeting of photospheric magnetic flux tubes (e.g., \\citealt{Roberts1979,Hollweg1981,Spruit1981, Choudhuri+others1993,Huang+others1995,Musielak+Ulmschneider2001}) or by small-scale reconnection \\citep{Axford+McKenzie1992}. Differentiating observationally between kink and Alfv\\'en wave modes is not entirely straightforward since both can cause transversal displacements and observationally appear as incompressible, at least for some magnetic field geometries \\citep{VanDoorsselaere+others2008}. In general, these wave modes can be distinguished in straight, cylindrical, high density magnetic flux tubes. Flux tubes act as wave guides for kink waves (i.e., the wave is confined in the tube), which in turn cause transverse displacements of the tube axis. In contrast, Alfv\\'en waves propagating incompressibly along a magnetic flux tube are torsional and do not lead to a transversal displacement \\citep{Erdelyi+Fedun2007,VanDoorsselaere+others2008}. Chromospheric kink and Alfv\\'en waves have been observed in off-limb spicules. Since fibrils are the likely on-disk counterpart of spicules (e.g., \\citealt{Tsiropoula+others1994}) one would expect to see transverse oscillations in them as well.     The observed properties of MHD waves can be used to measure the magnetic field strength in, e.g., coronal loops (e.g., \\citealt{Nakariakov+Ofman2001}) as well as spicules (e.g., \\citealt{Zaqarashvili+others2007,Singh+Dwivedi2007,Kim+others2008}). The thus measured spicule field strengths outside active regions are generally in the 10-40 G range. Spectropolarimetric measurements of spicules yield similar field strengths, e.g., 10 G from Hanle measurements of the \\ion{He}{1} 10830 \\AA\\ triplet \\citep{TrujilloBueno+others2005} and 30 G from \\ion{He}{1} D3 measurements \\citep{LopezAriste+Casini2005}. In this paper we present observations of transverse oscillations in on-disk fibrils, and based on them, give an estimate of the magnetic field strength in chromospheric active region dynamic fibrils. ",
        "conclusions": "We have presented on-disk observations of a kink wave in an active region chromospheric fibril. The high spatial and temporal resolution of the data have allowed us to study the event in detail. It begins with a lateral movement of the fibril which has its footpoint at the edge of a sunspot umbra.  The transverse oscillations of the fibril are associated with red-shifted, broadened intensity profiles, while indications of \\added{steepening waves} are seen in the footpoint. The period of the waves in the footpoint is roughly twice the period of the FOI transverse oscillations. The event ends as the fibril is retracted into the footpoint. It is not clear what triggers the event. One possible scenario is that the magnetic field above (at the base of the corona) changes and both the lateral movement and oscillations result from that. Since we see no \\added{strong} systematic differences in measured phases in our observations, we cannot say \\added{for sure} whether the disturbance propagates or, if it does, in which direction. \\added{Where a trend in measured phases is seen the propagation is downwards, i.e., towards the FOI footpoint.} Based on a comparison of Hinode observations and Monte Carlo simulations of Alfv\\'en waves in spicules, both upward and downward propagating waves are likely \\citep{DePontieu+others2007}.  If the observed event is due to reconnection higher up, the observed red-shift would be consistent with a bipolar jet occurring at the reconnection site (e.g., \\citealt{Innes+others1997}). The largest difference between the FOI intensity profile and the average intensity profile of the surroundings is at $\\approx$ 250 m\\AA\\, corresponding to a Doppler velocity of 9 km/s (Fig.~\\ref{fig:velstat}). Compared to velocities of -20 km/s seen in the \\ion{Ca}{2} 8542 \\AA\\ on-disk counterparts of type II spicules (rapid blue-shifted events, \\citealt{RouppevanderVoort+others2009}), the FOI velocity does not seem very high. \\cite{Shoji+Kurokawa1995} observed downflows of up to 50-100 km/s in H-$\\alpha$ and Ca K during the impulsive phases of two flares. They also observed \\ion{Na}{1} D1 and D2 lines as well as metallic lines which showed significantly smaller, up to 2-6 km/s, downflows. Based on the above, the magnitude of the FOI red-shift does not exclude a small-scale reconnection event higher up in the atmosphere as the trigger of the observed dynamics (no flares were recorded during the observation day, so it can only be a small event). A scenario involving a small scale reconnection event is also supported by the apparent change in the \\ca\\ Stokes $V$ polarity in the footpoint and the increased line and continuum intensities when the FOI connect to the footpoint. These can be interpreted as an emission feature near the line core caused by transient heating in the chromosphere \\citep{SanchezAlmeida1997}.   Fast and slow MHD waves may be coupled due to effects of boundary conditions, inhomogeneities, and gas pressure. Coupling of MHD waves is a possible explanation for the combined EIT and TRACE observations of loop footpoint intensity variations having the same period as the co-occuring fast transverse loop oscillation \\citep{Verwichte+others2010}. This was demonstrated by \\cite{Terradas+others2011}, who show how transverse motions of a coronal loop can produce slow waves with the same periodicity as the transverse oscillation. A coupling of MHD waves may be the connection between the observed FOI oscillations and the \\added{steepening} waves. Note that unlike in the coronal loop case, the observed \\added{steepening (indicating the presence of shocks)} wave has a period roughly twice that of the FOI oscillation.  The mean period, amplitude, and phase speed of the oscillations are similar to measurements in off-limb spicules (e.g., \\citealt{He+others2009} at heights below 2 Mm). The properties of the oscillations allowed us to estimate the field strength in the dynamic fibril to be at least $220 \\mathrm{\\mbox{\\,--\\,}} 330 \\mathrm{\\ G}$. For comparison, a value of 120 G at the base of the corona was used by \\cite{DePontieu+others2004} to model spicules via leakage of P-modes into the chromosphere. \\cite{Centeno+others2010} estimate, based on \\ion{He}{1} 10830 \\AA\\ triplet Hanle and Zeeman measurements, 50 G as a possible lower value for the field in network spicules. \\cite{LopezAriste+Casini2005} measured 30 G fields at a height of $3000 \\mathrm{\\mbox{\\,--\\,}} 5000 \\mathrm{\\ km}$ above the limb. Considering that we observed an active region and the magnetic field strength decreases with height (the \\ca\\ on-disk observations are at heights below 1500 km) a field strength of a few hundred Gauss is reasonable. It is also consistent with the chromospheric field strength near the outer boundary of a sunspot (e.g., \\citealt{ Solanki+others2003,OrozcoSuarez+others2005}). Once a telescope model for the \\ca\\ line becomes available, it will be interesting to compare this value with results from Zeeman measurements.  How common kink waves in fibrils are is difficult to estimate, though the similarity with kink waves in spicules suggests that the two should have similar occurrence rates. We see indications of a few more such oscillations in the data, but they are not as clear as the one presented here. For the oscillation to be visible the spatial and temporal resolutions need to be high. Also, the choice of wavelength may play a role: perhaps the oscillations were so well visible because the fibril was strongly red-shifted with respect to the surroundings. Due to the complex optically thick background consisting of other fibrils, the visibility of the oscillations on the disk is limited compared to the optically thin limb spicules."
    },
    "1107/1107.5115_arXiv.txt": {
        "abstract": "We present the first weak-lensing-based scaling relation between galaxy cluster mass, \\Mwl, and integrated Compton parameter \\Ys. Observations of 18 galaxy clusters at $z\\simeq0.2$ were obtained with the Subaru 8.2-m telescope and the Sunyaev-Zel'dovich Array. The \\my\\ scaling relations, measured at $\\Delta$~=~500, 1000, and 2500~$\\rho_c$, are consistent in slope and normalization with previous results derived under the assumption of hydrostatic equilibrium (HSE).  We find an intrinsic scatter in \\Mwl\\ at fixed \\Ys\\ of 20\\%, larger than both previous measurements of $M_\\mathrm{HSE}-\\Ys$ scatter as well as the scatter in true mass at fixed \\Ys\\ found in simulations.  Moreover, the scatter in our lensing-based scaling relations is morphology dependent, with 30 -- 40\\% larger \\Mwl\\ for undisturbed compared to disturbed clusters at the same \\Ys\\ at \\rfh. Further examination suggests that the segregation may be explained by the inability of our spherical lens models to faithfully describe the three-dimensional structure of the clusters, in particular, the structure along the line-of-sight.  We find that the ellipticity of the brightest cluster galaxy, a proxy for halo orientation, correlates well with the offset in mass from the mean scaling relation, which supports this picture. This provides empirical evidence that line-of-sight projection effects are an important systematic uncertainty in lensing-based scaling relations. ",
        "introduction": "Galaxy clusters are tracers of the highest peaks in the matter density field, and as such, their abundance as a function of mass and redshift depends strongly upon cosmology.  Along with data from a variety of other techniques \\citep{KomatsuE11,HickenE09,KesslerE09,PercivalE10}, galaxy cluster surveys have placed precise constraints on the $\\Lambda$CDM cosmological parameters \\citep{VikhlininE09-CCCP3,MantzE10}.  To improve the constraints from cluster surveys, the scaling relationship between survey observables and the total cluster mass must be precisely determined.  Surveys at millimeter wavelengths can be used to identify large numbers of clusters via the Sunyaev-Zel'dovich (SZ) effect, as has been demonstrated by the South Pole Telescope (SPT; \\citealt{StaniszewskiE09,VanderlindeE10,WilliamsonE11}), the Atacama Cosmology Telescope (ACT; \\citealt{MarriageE10}), and the {\\it Planck} satellite \\citep{planck-esz}.  The SZ effect is a spectral distortion of the cosmic microwave background (CMB) resulting from the inverse Compton scattering of CMB photons by the hot electrons of the intra-cluster medium (ICM).  The magnitude of this distortion is determined by the integral of the pressure along the line of sight through the cluster. The volume integral of pressure is the total thermal energy content of the ICM, which should directly trace the depth of the potential well \\citep{CHR02}, so a tight scaling between the SZ signal ($Y$, see Section~\\ref{sec:szmodel}) and mass is expected.  A low-scatter correlation is indeed found in simulations of galaxy clusters \\citep[e.g.,][]{daSilvaE04,MotlE05,Nagai06,StanekE10}, and is suggested in the best comparisons of SZ observations and cluster mass estimates available to date \\citep{BonamenteE08,PlaggeE10,AnderssonE11}.  However, the mass-$Y$ scaling relation remains poorly determined from these observations, and the uncertainty in the relation is already limiting the cosmological constraints derived from small SZ-selected samples \\citep{VanderlindeE10,SehgalE11}.  Moreover, previous comparisons of $Y$ with mass have relied on mass estimates derived assuming hydrostatic equilibrium (HSE).  This assumption is certain to result in at least some bias, although the degree of bias is poorly known; observational constraints suggest $\\lesssim20$\\% within radii of interest \\citep{MahdaviE08,ZhangE10}, though some simulations have found larger biases \\citep[e.g.,][]{CavaliereE11}.  One promising technique for improving our knowledge of the mass-$Y$ relation is weak gravitational lensing \\citep[e.g.,][]{DahleE02,HoekstraE07,OkabeE10}.  Unlike other techniques, weak lensing directly probes the gravitational potential of the cluster, providing a way to test for mass biases in other mass estimators.  However, the weak lensing signal is sensitive to the total mass projected along the line-of-sight, and thus likely suffers projection-related errors in mass measurements \\citep[e.g.,][]{MetzlerE01,Hoekstra01,Hoekstra03}.  Several authors have calibrated X-ray observables against lensing mass estimates, typically finding consistency with (but with larger scatter than) X-ray-only scaling relations \\citep{SmithE05,BardeauE07,HoekstraE07,ZhangE07,ZhangE08,RykoffE08,OkabeE10b,HoekstraE11}. This technique has not yet been used to calibrate the SZ observable except in the cluster core \\citep{MarroneE09}. An important next step is to make such a comparison at the larger radii typically used in cosmological studies with galaxy clusters.  The calibration of mass-observable relations is one of the key goals of the Local Cluster Substructure Survey (LoCuSS\\footnote{\\url{http://www.sr.bham.ac.uk/locuss}}), a multi-wavelength survey of galaxy clusters at $0.15 < z < 0.3$ selected from the {\\it ROSAT} All-sky Survey \\citep{EbelingE98,EbelingE00,BohringerE04}.  The large luminosity range and lack of morphological selection in the sample ensure a morphologically diverse population for calibration studies. Weak lensing mass estimates are available for an initial sample of 29 LoCuSS clusters (\\citealt{OkabeE10}; hereafter \\citetalias{OkabeE10}). In this paper, we combine these estimates with SZ data from the Sunyaev-Zel'dovich Array (SZA) in order to examine the mass-$Y$ scaling relation and its scatter.  We also explore the possibility of biases in the mass estimates related to cluster morphology, which our data suggest is an important systematic.  Section~\\ref{sec:observations} describes the sample of clusters as well as the weak lensing and SZ observations.  In Section~\\ref{sec:analysis}, we discuss the analysis techniques used to derive the weak lensing mass estimates and determine the SZ observables.  Our scaling relation results are presented in Section~\\ref{sec:results} and our interpretation of the findings is discussed in Section~\\ref{sec:discussion}.  Our conclusions are reviewed in Section~\\ref{sec:conclusions}. We assume a flat $\\Lambda$CDM cosmology: $\\Omega_\\mathrm{M}=0.27$, $\\Omega_\\Lambda=0.73$, $H_0=73\\rm~km~s^{-1}~Mpc^{-1}$. ",
        "conclusions": ""
    },
    "1107/1107.2445_arXiv.txt": {
        "abstract": "NGC~2419 is a massive outer halo Galactic globular cluster whose stars have previously been shown to have somewhat peculiar abundance patterns.  We have observed seven luminous giants that are members of NGC~2419 with Keck/HIRES at reasonable SNR.  One of these giants is very peculiar, with an extremely low [Mg/Fe] and high [K/Fe] but normal abundances of most other elements.  The abundance pattern does not match the nucleosynthetic yields of any supernova model.  The other six stars show abundance ratios typical of inner halo Galactic globular clusters, represented here by a sample of giants in the nearby globular cluster M30. Although our measurements show that NGC~2419 is unusual in some respects, its bulk properties do not provide compelling evidence for a difference between inner and outer halo globular clusters. ",
        "introduction": "}   NGC~2419 is a globular cluster (GC) in the outer halo of the Galaxy at a distance of 84~kpc\\footnote{The distance and other parameters for NGC~2419 are from the 2010 version of the on-line GC database of \\cite{harris96}, based on the photometry of \\cite{harris97}.}. \\cite{sigmav09} determined a velocity dispersion for this cluster based on 40 stars which leads to $M/L = 2.05\\pm0.50~M_{\\odot}/L_{\\odot}$, a normal value for an old stellar system, with no evidence for  dark matter at the present time.  \\citet*{cls10} further asserted that this cluster could not have formed in a now-defunct dark matter halo.  \\cite{harris97} obtained a deep CMD of NGC~2419 with HST/WFPC2 and demonstrated that there is no detectable difference in age between it and M92, an ancient, well-studied inner halo GC of comparable metallicity. Furthermore, there is no evidence from CMDs of multiple stellar populations in NGC~2419.   However, some of the characteristics of NGC~2419 are anomalous for a GC.  It has the highest luminosity ($M_V \\sim -9.6$~mag) of any GC with Galactocentric radius $R > 20$~kpc, higher than all other GCs with $R > 15$~kpc with the exception of M54, which is believed to represent the nuclear region of the Sgr dSph galaxy now being accreted by the Milky Way \\citep{ibata95,sarajedini95}.  NGC~2419 also has an unusually large half-light radius ($r_h = 19$~pc) and core radius ($r_c = 7$~pc) for a massive Galactic GC and a relaxation time which exceeds the Hubble time, also unusual for a GC.  Every other luminous GC in the outer halo is considerably more compact.  The 2010 version of the on-line database of \\cite{harris96} lists only three Galactic GCs that lie anomalously above the bulk of the Galactic GCs in the plane $r_h$ vs $M_V$, one of which is NGC~2419. The other two are M54 and $\\omega$ Cen, the most luminous Galactic GC, and one with a large internal spread in metallicity, also believed to be the stripped core of a dwarf galaxy accreted by the Milky Way.  For these reasons, \\cite{vdb04} and \\cite{mackey05}, among others, suggested that NGC~2419 is also the remnant of an accreted dwarf galaxy.  In fact, \\citet{irwin99} and \\citet{newberg03} suggested that NGC~2419 is one of many GCs associated with the disrupted Sgr dSph, but \\citet{law10} firmly concluded that it is not.  Therefore, the cluster is potentially a piece of an unidentified, tidally disrupted dSph.    To follow up on these earlier suggestions, \\cite{deimos} obtained moderate resolution spectra of a large sample of luminous giants in NGC~2419 and analyzed line strengths in the region of the near-infrared Ca triplet.  We found that there is a small but measurable spread in Ca abundance of $\\sim$0.3~dex within this massive metal-poor globular cluster.  We present here an analysis of high dispersion spectra of seven luminous giant members of NGC~2419, whose chemical inventory we compare to that inferred from a similar sized sample of luminous RGB members of the nearby inner halo globular cluster M30 (NGC~7099), observed and analyzed in the same way as the NGC~2419 stars. Note that these two globular clusters have the same metallicity in the 2003 version of the on-line database of \\cite{harris96}. We explore two issues:  \\begin{enumerate}  \\item Is the chemical inventory in such a distant cluster identical to that of a globular cluster of similar [Fe/H] in the inner halo?  \\item What star-to-star variations are revealed by our high-dispersion spectra?   \\end{enumerate} ",
        "conclusions": "Our key result from the analysis of 7 luminous giant members of the distant outer halo globular cluster NGC~2419 is that one of these stars is extremely peculiar, having a very low [Mg/Fe] ratio, but being normal in all other element ratios except for a high [K/Fe]. Similarly low [Mg/Fe] ratios are seen in stars in dSph satellites of the Galaxy, specifically in Ursa Minor \\citep{cohen_umi} and in Sextans \\citep{shetrone01}.  But in both cases these stars have low values of other $\\alpha$ elements, which is not the case for NGC~2419 S1131. J.~Cohen has examined and analyzed spectra of $\\gtrsim$100 stars in Galactic GCs, and has never seen one similar to S1131 in NGC~2419.  Furthermore, the HIRES spectra support the small range in metallicity within the giants in NGC~2419 found by \\cite{deimos}, which might suggest that this distant outer halo GC is the remnant of an accreted dwarf galaxy. A definitive result will require further observations of the faint upper RGB stars in this GC, concentrating on those whose moderate resolution Deimos spectra from \\cite{deimos} suggest they are more metal-rich than the bulk of the stars in this cluster, which will be carried out in  a future campaign.   Ignoring the limited peculiarities of S1131, we find that there is no substantive difference between the mean behavior of the abundance ratios for six other red giants in NGC~2419 and those ratios characteristic of the inner halo despite the structural peculiarities of NGC~2419. The influence of star formation rates and other aspects of the detailed chemical evolution of a stellar system only affect such ratios once  star formation has been underway long enough for sources other than SNII to become important contributors, which does not seem to be the case in very metal-poor systems such as NGC~2419, nor in the most metal-poor stars in the dSph satellites of the Milky Way (Cohen \\& Huang 2009, 2010; Kirby et al 2011).  Thus, like the previous studies based on detailed abundance analyses of individual or co-added stellar spectra of stars in distant outer halo clusters, the chemical inventory of the globular cluster system appears to be independent of location.  With the exception of those clusters known to be associated with the currently disrupting Sgr dSph galaxy, every globular cluster exhibits the same trends as a function of overall cluster metallicity---parameterized by [Fe/H]---irrespective of kinematics or location within the halo.  The existence of a metallicity gradient in the Galactic halo is controversial.  \\citet{carollo07} found evidence for a break in the metallicity distribution of halo field stars in the sense that the more distant stars are more metal-poor;  this conclusion is supported by \\cite{rix10}.  Furthermore, \\cite{carollo07} showed that the metallicity within the outer halo population displays a gradient toward even lower metallicities for the most distant stars or stars showing the highest magnitude of rotation retrograde to the Milky Way disk.  However, both \\citet{sesar11} and \\citet{schoenrich10} improved upon \\citeauthor{carollo07}'s study.  These new works, which are disputed by \\cite{beers11}, found no evidence for a dichotomous halo, particularly in regard to the kinematics and density profile.  Our measurements for NGC~2419 additionally provide evidence against a dichotomy in detailed abundances.  The outer halo globular clusters show the same bulk abundance properties---and, by extension, (short) star formation timescales---as the inner halo globular clusters."
    },
    "1107/1107.5579_arXiv.txt": {
        "abstract": "We present hydrodynamical simulations of major mergers of galaxies and study the effects of winds produced by active galactic nuclei (AGN) on interstellar gas in the AGN's host galaxy. We consider winds with initial velocities $\\sim 10,000 \\kms$ and an initial momentum (energy) flux  of $\\sim \\tau_w L /c$  ($\\sim 0.01 \\, \\tau_w \\, L$), with $\\tau_w \\sim 1-10$.   The AGN wind sweeps up and shock heats  the surrounding interstellar gas, leading to a galaxy-scale outflow with velocities $\\sim 1000 \\kms$, peak mass outflow rates comparable to the star formation rate, and a total ejected gas mass $\\sim 3 \\times 10^9 \\, M_\\odot$.  Large momentum fluxes, $\\tau_w \\gtrsim 3$, are required for the AGN-driven galactic outflow to suppress star formation and  accretion in the black hole's host galaxy.  Less powerful AGN winds ($\\tau_w \\lesssim 3$) still produce a modest galaxy-scale outflow, but the outflow has little global effect on the ambient interstellar gas. We argue that this mechanism of AGN feedback can plausibly produce the high velocity outflows observed in post-starburst galaxies and the massive molecular and atomic outflows observed in local ultra-luminous infrared galaxies.   Moreover, the outflows from local ultra-luminous infrared galaxies are inferred to have  $\\tau_w \\sim 10$, comparable to what we find is required for AGN winds to regulate the growth of black holes and set the $M_{BH}-\\sigma$ relation.   We conclude by discussing theoretical mechanisms that can lead to AGN wind mass-loading and momentum/energy fluxes large enough to have a significant impact on galaxy formation. ",
        "introduction": "\\label{sec:intro}  Accretion onto a central massive black hole (BH) in a galactic nucleus produces energy in the form of radiation, relativistic jets, and wider angle (less-collimated) non-relativistic ($v \\sim 10^4 \\kms$) outflows \\citep{krolik99}.   The coupling of this energy output to gas in galaxies and in the intergalactic medium is believed to play an important role in galaxy formation, potentially regulating the growth of massive galaxies and the thermal properties of the intracluster medium in galaxy groups and clusters (e.g., \\citealt{silk98,croton06}).  The impact of this `feedback' on the gas {\\em in} galaxies is particularly uncertain, both because the interstellar (ISM) gas is denser, and thus more difficult to affect dynamically, and because much of the ISM subtends a relatively modest solid angle relative to a central active galactic nucleus (AGN).   However, analytic estimates and numerical simulations have demonstrated that if a modest fraction of the energy produced by accretion onto a central BH can couple to the surrounding gas, it can unbind the interstellar gas (e.g., \\citealt{silk98,dimatteo05}).  The physical processes most likely to produce such an effect are winds \\citep{king03, king11}, radiation pressure \\citep{murray05}, and/or Compton heating \\citep{sazonov04} from a central AGN.   Understanding how this works in detail is one of the major challenges in our understanding of the connection between AGN physics and galaxy formation.  In this paper, we assess the influence of AGN winds on gas in the AGN's host galaxy using three-dimensional numerical  simulations.   Previous analytic work and one and two-dimensional simulations have demonstrated that  AGN winds can in principle sweep up and drive gas out of galaxies, potentially explaining the $M_{BH}-\\sigma$ relation and the dearth of gas and ongoing star formation in massive, early-type, galaxies (e.g., \\citealt{king05,king11,novak10,ostriker10} and references therein).  Observationally, there is strong evidence that AGN indeed drive powerful outflows.    Broad-absorption line (BAL) quasars, which show blue-shifted absorption lines in the rest-frame ultraviolet with inferred outflow velocities $\\sim 10,000-40,000 \\kms$, represent over $\\sim 40\\%$ of quasars in infrared selected samples \\citep{dai08}. A similar fraction of radio-quiet quasars show evidence for high velocity outflows in X-ray absorption line spectroscopy  \\citep{tombesi10}.  It is likely that {\\em all} quasars possess such outflows but that they are only observed when the system is viewed modestly edge-on \\citep{murray1995}.   However, determining the mass-loss rate from spatially unresolved absorption-line observations is notoriously difficult given uncertainties in the radius of the absorbing gas. In a handful of low-ionization BAL quasars (in particular, FeLoBALs) this degeneracy has been broken, suggesting mass loss rates significantly larger than the black hole accretion rate \\citep{moe09, bautista10,dunn10,claude11}.   These observations trace absorbers at large distances from the BH ($\\sim$ kpc), in contrast to most of the high ionization UV and X-ray absorption seen in BAL quasars, which arises at $\\lesssim 1$ pc.   In addition to these well-characterized outflows, it is possible that quasars drive even more powerful outflows during phases of super-Eddington accretion \\citep{kingpounds03, king11} or during obscured phases when the AGN's radiation  is trapped in the galactic nucleus, enhancing the force on the ambient gas \\citep{debuhr10, debuhr11}.  To characterize the impact of AGN winds on their host galaxy, we carry out numerical simulations of major galaxy mergers with models for BH growth and feedback.   Although it is by no means certain that all quasars are associated with mergers, this provides a convenient and well-posed model in which to study gas inflow in galactic nuclei.  It also allows us to readily compare our results to the extensive previous literature on AGN feedback during mergers.  In the next section (\\S \\ref{sec:methods}) we summarize our methods, emphasizing our implementation of AGN winds and our treatment of ISM cooling.   We then describe our key results in \\S \\ref{sec:results}. We conclude in \\S \\ref{sec:discussion} by discussing the implications of our results for models of AGN feedback and for galactic winds driven by AGN.   We also compare our results to observations of high-velocity outflows from post-starburst galaxies and ultra-luminous infrared galaxies and summarize the theoretical processes that can produce the large mass-loading and momentum fluxes we find are necessary for AGN winds to have a substantial impact on the surrounding ISM. ",
        "conclusions": "\\label{sec:discussion}  We have carried out three-dimensional SPH simulations of the interaction between a high-speed outflow produced by an AGN and interstellar gas in the AGN's host galaxy.  We show that this interaction can drive a large-scale galactic outflow that in some cases unbinds a significant fraction of the ISM of the host galaxy (see, e.g., \\citealt{king05} for closely related arguments).  The AGN-driven galactic winds found here provide a possible explanation for the high velocity outflows observed in some post-starburst galaxies \\citep{tremonti07} and for the massive atomic and molecular outflows with $v \\sim 1,000 \\kms$ seen in local ultra-luminous infrared galaxies (ULIRGs) \\citep{feruglio10,chung11,rupke11,sturm11}.   We return to this below.  Our specific calculations assume that AGN activity is triggered by major mergers of galaxies (as seen, for example, by \\cite{koss2010}), but we suspect that the results presented here are  much more general and apply relatively independent of the physical mechanism(s) that drive gas into the central $\\sim 100$ pc of galaxies.   The physically-motivated model utilized here for AGN wind feedback is easy to implement (\\S \\ref{sec:windfeedback}) and could readily be applied in other contexts, including high resolution simulations of galactic nuclei and cosmological simulations of BH growth and evolution.  Our results provide a quantitative mapping between the properties of an AGN wind in the vicinity of the BH and the resulting large-scale galaxy-wide outflow (Table \\ref{tab:simparm} and Fig.~\\ref{fig:woblowout}).  We have parametrized the AGN wind at small radii in terms of its speed $v_w$ and momentum flux $\\dot p_w = \\tau_w L/c$ -- the corresponding energy flux in the wind is $\\dot E_w = 0.5 \\, \\dot p_w \\, v_w = 0.5 \\, \\tau_w (v_w/c) \\, L$.  We find that an AGN wind with $v_w \\sim 10,000 \\kms$ -- typical of, e.g., broad-absorption line quasars -- produces a galactic outflow with velocities $\\sim 1,000 \\kms$.    This reduction in velocity is a consequence of the AGN wind sweeping up and driving out $\\sim 5-10$ times as much mass as was in the wind initially (e.g., Fig.~\\ref{fig:vpdf}).   Quantitatively, we find that the AGN-driven galactic outflow unbinds $\\sim 1-5 \\times 10^9 M_\\odot$ of gas, $\\sim 10-40 \\%$ of the initial gas in the two merging galaxies (which are Milky Way-like spirals in our calculations).    This is true for a wide range of outflow models, covering $\\tau_w \\sim 1-10$, $v_w \\sim 7,000-10,000 \\, \\kms$, and with and without the radiation pressure feedback model from \\citet{debuhr10,debuhr11}.   Not surprisingly, AGN winds with larger momentum and energy fluxes (larger $\\tau_w$ and $v_w$) are more effective at driving galaxy-scale outflows (Fig.~\\ref{fig:woblowout}).  The mass outflow rate from the vicinity of the galaxy is somewhat noisy and thus difficult to directly measure in the simulations.  The average outflow rate relative to the star formation rate can, however, be readily estimated by comparing the total unbound mass ($M_{\\rm out}$ in Table \\ref{tab:simparm}) with the mass in new stars formed in the simulation ($M_{*,f}$ and $M_{*,p}$ in Table \\ref{tab:simparm}, where $M_{*,f}$ is the total mass of new stars formed and $M_{*,p}$ is the mass of new stars formed after the peak of star formation).   This comparison implies that the average mass outflow rate is typically comparable to the star formation rate; after the peak of star formation in the merger, which is when the BH reaches its final mass, the average outflow rate often exceeds the corresponding star formation rate by factors of a few.  Physically, we find that the properties of the galaxy-wide outflow are determined by both the momentum and energy fluxes in the AGN wind at small radii.   A larger momentum flux implies that more mass can be swept out of the galaxy while a larger energy flux in the wind (e.g., larger $v_w$ at fixed $\\dot p_w$) leads to more shock heating of the ambient ISM.   This shocked gas partially cools but retains enough thermal energy to contribute to driving the galaxy-wide outflow.   For quantitatively accurate results we find that it is necessary to modify the \\citet{springel03} equation of state to properly include metal-line cooling and Compton heating/cooling as processes by which gas relaxes back to the effective equation of state that is used to calculate the ISM pressure (see \\S \\ref{sec:cooling}).  To produce a galaxy at the end of the simulation that is on the $M_{BH}-\\sigma$ relation using only feedback by our AGN wind model, we require $\\tau_w \\gtrsim 5-10$ (see Table~\\ref{tab:simparm} and Fig.~\\ref{fig:woblowout}; the velocity dispersion of the remnant galaxy is $\\simeq 170 \\kms$ so that $M_{BH} \\simeq 7 \\times 10^7 M_\\odot$ for galaxies on the mean $M_{BH}-\\sigma$ relation).   This required momentum flux is comparable to, although somewhat smaller than, the required momentum flux of $\\dot p \\sim 20 \\, L/c$ that DQM found  using a simple model of radiation pressure feedback in which all particles in the vicinity of the BH feel a force $\\propto \\tau \\, L/c$.   The energy flux corresponding to $\\tau_w \\sim 5-10$ and $v_w \\sim 10,000 \\kms$ is $\\dot E \\sim 0.1 \\, L$, similar to the energy injection rate of $\\sim 0.05 \\, L$ that \\citet{dimatteo05} found was required to reproduce the $M_{BH}-\\sigma$ relation (although they deposited thermal energy into ambient gas while we supply kinetic energy to a wind).  One  important distinction that we find is between AGN driving galactic outflows and AGN outflows regulating the growth of the BH.   For example, in our models with $\\tau_w = 1$ and $v_w = 10,000 \\kms$, the final BH mass is a factor of $\\sim 10$ above the mean $M_{BH}-\\sigma$ relation in calculations that only include AGN wind feedback.   Also including  the simple radiation pressure feedback model from DQM in the simulation resolves this discrepancy; the same AGN wind model then unbinds $\\sim 10^9 M_\\odot$ of gas, $\\sim 10 \\%$ of the initial gas mass in the system.  However, in spite of this outflow, the AGN wind neither regulates the BH growth nor has a significant effect on the average star formation in the system.   This demonstrates that very powerful AGN outflows with $\\tau_w \\sim 3-10$ are required for the BH to have a substantial effect on either its own accretion history or the star formation in the surrounding galaxy. This conclusion is similar to that of \\cite{silk10}, who required a comparable momentum to eject gas from their model galaxies, and to that of \\cite{kaviraj11}, who found that simple models for the color evolution of early type galaxies required feedback similar in strength ($\\dot{E}$ of a few percent) and duration ($\\sim 200$ Myr) to what we find here.  The galactic outflows found in our simulations with $v_w \\sim 10,000 \\kms$ and $\\tau_w \\gtrsim 3$ have properties broadly similar to those observed in some massive post-starburst galaxies \\citep{tremonti07} and in local ULIRGs \\citep{feruglio10,chung11,rupke11,sturm11}.   Moreover, it is striking that \\citet{sturm11} inferred $\\dot p_w \\sim 10  \\, L/c$, i.e., $\\tau_w \\sim 10$, for the molecular outflows in local ULIRGs (these values are uncertain at the factor of $\\sim 3$ level; see, e.g., \\citealt{sturm11} for more details).  This is comparable to the outflow momentum flux we find is required to {both} suppress late time star formation during galaxy mergers and produce remnants approximately on the observed $M_{BH}-\\sigma$ relation. This correspondence is encouraging, though it does not directly address the question of what mechanism powers such outflows in the first place.   The relative role of star formation and AGN in  powering the $\\sim 1000 \\kms$ molecular outflows is also  unclear, with \\citet{sturm11} presenting evidence from OH observations in favor of AGN, while \\citet{chung11} argue that star-formation is the dominant power source based on their CO observations.  In addition, in some samples of post-starburst galaxies, the observed outflows appear consistent with those driven by star formation \\citep{coil11}.  From the theoretical perspective, the biggest uncertainty in applying our results is the uncertainty in the absolute momentum/energy flux in AGN outflows at small radii.   This uncertainty remains even for the most well-characterized outflows, those of broad-absorption line (BAL) quasars.  In theoretical models in which BAL outflows are produced by radiation pressure on lines in the vicinity of the broad-line region, the predicted momentum flux is typically  $\\lesssim  L/c$, i.e., $\\tau_w \\lesssim 1$ \\citep{murray1995}.  This follows  from the fact that (observationally) the lines  do not cover the entire continuum and thus cannot absorb the entire continuum momentum flux.   This constraint is less clear for low-ionization BALs (in which the lines cover more of the continuum), but probably remains reasonably accurate.   One way out of this conclusion is if magnetic fields contribute to accelerating BAL outflows (as in, e.g., \\citealt{proga03}) since this can increase both the mass flux and terminal velocity of the outflow.    There are observational estimates that the momentum flux in some low-ionization BAL outflows may indeed be $\\gtrsim L/c$  \\citep{moe09, bautista10,dunn10}, though these are difficult observations to interpret and they have complicated selection effects.  In addition to line-driven winds, it is possible that even more powerful outflows are driven during evolutionary phases distinct from optically bright quasars.   The two most promising from our point of view are:  (1) outflows driven by radiation pressure on dust during phases when the ISM around the AGN is optically thick to far-infrared radiation \\citep{murray05,debuhr10}.   In this case the momentum flux in a small-scale AGN wind can in principle reach $\\sim \\tau L/c$, where $\\tau$ is the infrared optical depth of the galactic nucleus.  Such an outflow is not directly seen in DQM's simple implementation of radiation pressure feedback.  However, this could easily be a limitation of the lack of a proper treatment of the infrared radiative transfer.  Indeed, we expect that radiation pressure will inevitably drive a high-speed outflow off of the nuclear disk (`torus') at $\\sim 1-10$ pc.    (2)   An additional source of powerful high-speed outflows can be produced if the fueling rate of the central AGN significantly exceeds Eddington (e.g., \\citealt{king03}).  In this case, the accretion disk becomes radiatively inefficient and much of the nominally inflowing mass is likely unbound (e.g., \\citealt{blandford99}).   The disparity in spatial scales between the horizon and where the AGN accretion rate is set ($\\sim 1-10$ pc; e.g., \\citealt{hopkins10c}) implies that it is very likely that there is a phase of super-Eddington fueling.  The critical question is whether this phase lasts long enough to dominate the {integrated} properties of a black hole's outflow.  All of the calculations presented in this paper utilize the effective equation of state model of \\citet{springel03}  (modified as in \\S \\ref{sec:cooling}) to pressurize the ISM and prevent runaway gravitational fragmentation.   This is a very simplistic treatment of ISM physics and in the future it is important to study the interaction between a central AGN and the surrounding ISM using more realistic  models that approximate the turbulent, multi-phase structure of the ISM (e.g., \\citealt{dobbs11,hopkins11})."
    },
    "1107/1107.2390_arXiv.txt": {
        "abstract": "Our new formulation of the Spherical Collapse Model (SCM-L) takes into account the presence of angular momentum associated with the motion of galaxy groups infalling towards the centre of galaxy clusters. The angular momentum is responsible for an additional term in the dynamical equation which is useful to describe the evolution of the clusters in the non-equilibrium region which is investigated in the present paper. Our SCM-L can be used to predict the profiles of several strategic dynamical quantities as the radial and tangential velocities of member galaxies, and the total cluster mass.  A good understanding of the non-equilibrium region is important since it is the natural scenario where to study the infall in galaxy clusters and the accretion phenomena present in these objects. Our results corroborate previous estimates and are in very good agreement with the analysis of recent observations and of simulated clusters. ",
        "introduction": "\\label{sec:intro}   The cosmic structures we observe in the Universe are thought to be originated by small perturbations in the primordial  matter distribution, enhanced in time by self-gravity. Density perturbations decouple from the Hubble flow and drive the collapse of matter into structures (groups and clusters of galaxies) through a highly non-linear process. In a Universe with cold dark matter and a cosmological constant ($\\Lambda$-CDM), structures of smaller scale decouple before structures of larger scales do: this means that galaxy clusters are built up through merging and coalescence of smaller galaxy groups. This ``bottom-up'' scenario has been widely tested using large $N$-body hydrodinamical simulations and is consistent with the most recent measurement of the cosmological parameters.  An accurate general description of the cosmic structure collapse is analitically inpraticable. To tackle the problem, simplifying assumptions are required. The most common and accepted assumption is to consider only spherically symmetric density perturbations: in this case, each perturbation can be described as an independent Friedmann universes on its own, with constant positive curvature, which first expands and then collapses mantaining its spherical shape. This Spherical Collapse Model (SCM) was first developed by \\cite{GG}, \\cite{S74}, and \\cite{G}, and it has since became a widely accepted approach in the literature (e.g.~\\citealt{P76,P80,Sch,Shm,Oal}). Many authors demonstrated that despite its simplistic assumption the SCM provides a quite good description of the dynamics in the outskirts of clusters, where the matter is not in energy equilibrium and is characterized by an overall infall motion towards the cluster centre (e.g.~\\citealt{Pal,BRal,BC,CMM08}, hereafter CMM08). Corrections to the SCM are generally required in the cluster cores, due to the interpenetration of different collapsing shells of matter (an effect referred to as shell crossing; see e.g.~\\citealt{SC} and references therein).  One of the major shortcomings of the SCM is that in principle a purely spherically symmetric scenario does not allow for any non-radial motion of the collapsing matter. In fact, galaxy dynamics in observed clusters show a significant amount of tangential motion \\citep{RG}. \\cite{DG} demonstrated that the amplitude of the caustic surfaces observed in the redshift-space distribution of galaxies is related to the galaxy escape velocities, which are by definition non-radial (see e.g. \\citealt{D99,Ral00,Ral01,Ral03,RD}; and \\citealt{D09} for a review). More recently, \\cite{CMM10} showed that radial velocities and tangential velocities of galaxies are statistically of the same order of magnitude in modulus. Observation of line-of-sight velocities can therefore be used to estimate the infall pattern and the mass profile of clusters up to the outermost regions (corresponding to radii exceeding seven times the virialization radius).  In this paper we take into account the effect of tangential motion into the SCM. Considerations based on the $\\Lambda$-CDM bottom-up scenario suggest that angular momentum in cluster outskirts is mainly associated with galaxy groups which belongs to the clusters, rather than single galaxies. In fact, galaxies rarely infall isolated, since groups appear to be formed long before they merge into clusters. We will show that the tangential motion of infalling galaxies is not only due to the intrinsic velocity dispersion of the infalling groups which the galaxies belong to, but also to the coupled effects of the angular momentum conservation and the torque produced by the presence of nearby structures. Several authors discussed the effect of the angular momentum in cosmic structures, but they mainly focused on the growth of the total angular momentum and/or on the behaviour of the inner, relaxed regions (e.g.~\\citealt{RG,R88,CT,ARal,N,Aal,DPK}). On the contrary, we will not take into account here the overall angular momentum of the galaxy cluster which can arise from the misalignment between the inertia tensor and the tidal forces acting on the collapsing structures \\citep{W}, nor the process of violent relaxation \\citep{LB}, which transforms the infall energy into non-radial velocity dispersion on a collapse timescale. But we will follow a different approach and focus on the non-equilibrium region of clusters, where the presence of angular momentum is expected to slow down the overall infall motion of matter. We will show that the SCM with angular momentum (hereafter SCM-L) better describes, compared to the standard SCM, the infall velocity profile and the mass profile in the outskirts of cluster; in fact it predicts shallower infall velocity profiles which are consistent with the results of CMM08 about the velocity profiles.  We wish to stress that considering the effects of the angular momentum in the dynamics of the outskirts of clusters necessarily implies an approach which is linked to the dynamics of the separate structures infalling towards the cluster core, so we have not to deal with the cluster total angular momentum.  In the following, we will discuss how to to take into account the effect of the angular momentum and how to derive it from the power spectrum of the density perturbations which originated the clusters (section \\ref{sec:proc}). The derived  equation provides information about both the matter distribution and the kinematical properties of clusters. We will thoroughly compare these predictions with our results published in CMM08 and CMM10 (section \\ref{sec:results}) and discuss their implications in cosmology (section \\ref{sec:concl}). The details of the calculation are generally omitted in the main text, and a full account on them can be found in the appendices.  In the present paper we adopt $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:concl}  Our new formulation of the Spherical Collapse Model (SCM-L) takes into account the presence of angular momentum associated with the motion of galaxy groups infalling towards the centre of galaxy clusters. The angular momentum (estimated from the power spectrum of density perturbations) is responsible for an additional term in the dynamical equation which is useful to describe the evolution of the clusters in the non-equilibrium region. Once the evolution equation is solved, our SCM-L provides a complete description of the dynamical state of the perturbations at any given redshift $z$ as a function of the adopted values of matter density function $\\Omega_0$.  We run our model with different values $\\Omega_0$ between $0.2$ and $0.4$ and redshift $z$ between $0$ and $0.5$. In each case, we computed the turnaround overdensity $\\Delta_\\rmn t$ and the ratio between the turnaround scale and the virialization scale.  Our SCM-L can be also used to predict the profiles of several strategic dynamical quantities (radial and tangential velocities of member galaxies, and total cluster mass). The agreement between our SCM-L and the simulations is always remarkable in the whole non-equilibrium region (spanning approximately from $0.5$ to $2$ turnaround radii $r_\\rmn t$), which is the region of interest of the present paper.  The non-equilibrium region is the natural scenario where to study the infall in galaxy clusters and the accretion phenomena present in these objects. The method described in the present paper can be also generalized to alternative cosmological models.  Our results corroborate our previous estimates (CMM08); moreover, they are in very good agreement with the analyses of recent observations \\citep{DG} and of simulated clusters (see, e.g., CMM10).%"
    },
    "1107/1107.2359_arXiv.txt": {
        "abstract": "{We present a new method to search for heavy nuclei sources, on top of background, in the Ultra-High Energy Cosmic Ray data~\\cite{Giacinti:2010ep}. We apply it to the 69 events with energies $E \\geq 55$\\,EeV published by the Pierre Auger Collaboration. We find a set of events for which the method reconstructs the source near the Virgo galaxy cluster. The probability to have a comparable set of events in some background is $\\sim 0.7$\\%. The reconstructed source is located at $\\simeq 8.5$ degrees from the active galaxy M87. The probability to reconstruct the source at less than 10 degrees from M87 for data already containing a comparable set of events is $\\sim 0.4$\\%. This may be a hint at the Virgo galaxy cluster as an ultra-high energy heavy nuclei source~\\cite{Semikoz:2010cc,Giacinti:2010ep}. We discuss the capability of current and near future experiments to test this possibility. Such a scenario gives a self-consistent description of the Auger anisotropy and composition data at the highest energies.} ",
        "introduction": "The Pierre Auger Observatory measurements of the Ultra-High Energy Cosmic Ray (UHECR) composition are compatible with a shift towards heavy nuclei, above $\\sim 10^{19}$\\,eV~\\cite{Collaboration:2010yv}. The analysis of the muon data from the Yakutsk EAS Array also hint at a significant fraction of heavy nuclei at the highest energies~\\cite{Glushkov:2007gd}. On the contrary, both HiRes measurements~\\cite{BelzICRC} and preliminary results of Telescope Array~\\cite{TA} are still consistent with a proton composition.  Both for protons and iron nuclei, UHECR sources must be located in the local Universe, within $r \\leq 150$\\,Mpc for energies $E>6 \\times 10^{19}$\\,eV. They should lie within the Large Scale Structure (LSS) of galaxy distribution.  Until now, methods to look for sources have been proposed for proton or light nuclei primaries~\\cite{Golup:2009cv,Giacinti:2009fy}. We present here a new method to search for UHE heavy nuclei sources. In most regions of the sky, one cannot detect heavy nuclei sources without a better knowledge of the Galactic Magnetic Field (GMF) than currently available~\\cite{Giacinti:2010dk,Giacinti:2011uj,ICRC_0229}. Nonetheless, we show that in some favourable cases, one may construct an algorithm to detect some of such sources with the present knowledge of the GMF.  We apply this method to the 69 Auger events published in Ref.~\\cite{Collaboration:2010zzj}. We detect a set of events for which the associated reconstructed source lies near Virgo. The Virgo cluster is in principle a good candidate for containing one or several UHECR source(s). The probability to have a comparable set of events in some background and reconstruct the source in any direction of the sky is $\\sim 7\\times10^{-3}$. If such a set of events already exists in the data, the probability to reconstruct the source at less than 10$^{\\circ}$ from M87 is $\\sim 4\\times10^{-3}$. The possibility that the detected feature is the heavy nuclei image of Virgo would be compatible with both the Auger composition and anisotropy at the highest energies. However, the present statistics do not allow us to conclude firmly. We discuss the ability of current and near future experiments to confirm or rule out this possibility.  We present the method in Section~\\ref{Method}. In Section~\\ref{Data}, we analyse the Auger data of Ref.~\\cite{Collaboration:2010zzj} and report the noteworthy feature. Section~\\ref{Discussion} presents a discussion of the results. ",
        "conclusions": "We have proposed here both a new method to detect some UHE heavy nuclei sources in the data, and a new interpretation of the Auger data which is both consistent with its composition and anisotropy measurements.  In Section~\\ref{Method}, we showed that some heavy nuclei sources located in some parts of the sky may have at least one of their images which looks like an enlarged proton source image. We presented a method which is able to detect such favourable sources on top of background. In Section~\\ref{Data}, we applied this method to the 69 Auger events with $E>55$\\,EeV of Ref.~\\cite{Collaboration:2010zzj} and found that the 142\\,EeV event and the Cen~A region events may be the UHE heavy nuclei image of the Virgo cluster. The associated reconstructed source is indeed located at only a few degrees from Virgo, at $\\simeq 8.5^{\\circ}$ from M87. The probability to reconstruct the source at less than $10^{\\circ}$ from M87 in some data already containing a comparable set of events is $\\sim 0.4$\\%. A better knowledge of the GMF than currently available, and more UHECR data are however still needed to confirm or rule out this possibility.  In Section~\\ref{Discussion}, we cross-checked this possibility with a ``blind-like'' analysis. We also pointed out that the anisotropy in the Auger data is due to the Cen~A region events. We note that the next generation of UHECR experiments will be able to test the possibility presented here."
    },
    "1107/1107.4263_arXiv.txt": {
        "abstract": "An algorithm is proposed for constructing a group (ensemble) pulsar time based on the application of optimal Wiener filters. This algorithm makes it possible to separate the contributions of variations of the atomic time scale and of the pulsar rotation to barycentric residual deviations of the pulse arrival times. The method is applied to observations of the pulsars PSR B1855+09 and PSR B1937+21, and is used to obtain corrections to UTC relative to the group pulsar time PT$_{\\rm ens}$.  Direct comparison of the terrestial time TT(BIPM06) and the group pulsar time PT$_{\\rm ens}$ shows that they disagree by no more than $0.4\\pm0.17\\; \\mu$s. Based on the fractional instability of the time difference TT(BIPM06) -- PT$_{\\rm ens}$, a new limit for the energy density of the gravitational-wave background is established at the level $\\Omega_g {h}^2\\sim 10^{-9}$. ",
        "introduction": "An observer located on the Earth, which is rotating about its axis and revolving around the Sun, receives a signal from a pulsar using a radio telescope over an accumulation time that is sufficient to obtain a specified signal-to-noise ratio. The pulse arrival time (PAT) is measured in a local time scale via crosscorrelation with a standard profile for the pulsar pulse.  The resulting PATs $\\tau_N$ can be translated into UTC, TAI, and TT using the relations \\cite{oleg90} \\begin{equation} {\\rm UTC}=\\tau_N+\\Delta\\tau,\\;{\\rm TAI}={\\rm UTC}+k,\\; {\\rm TT}={\\rm TAI}+32.184\\,{\\rm s}, \\end{equation} where  $\\Delta\\tau$ is the local time correction, $k$ is a whole number of seconds introduced to take into account the variation in the length of the day, and the quantity 32.184 s is added to remove a historical jump between atomic and ephemerides time. Since the duration of the second in TT depends on the position and velocity of the Earth in its orbit, an additional transformation from TT to Barycentric Time TB is required, using the algorithm described in \\cite{fairhead90}.  The observations that have been transformed into the uniform Barycentric Time system TB must then be reduced to the barycenter of the solar system, in order to exclude variations in the PATs due to the motion of the observer about the barycenter \\cite{oleg90} \\begin{equation} T=t-t_0+\\Delta_R(\\alpha,\\delta,\\mu_\\alpha,\\mu_\\delta,\\pi)+\\Delta_{\\rm orb}-DM/f^2+\\Delta_{\\rm rel}, \\end{equation} where $t_0$ is the initial epoch;  $t$ the topocentric PAT in TB; $T$ the PAT at the barycenter of the solar system; $\\Delta_R(\\alpha,\\delta,\\mu_\\alpha,\\mu_\\delta,\\pi)$ the Remer correction along the Earth's orbit; $\\alpha,\\delta,\\mu_\\alpha,\\mu_\\delta$ and $\\pi$ the right ascension, declination, proper motion in right ascension and declination, and parallax of the pulsar, respectively; $\\Delta_{\\rm orb}$ the Remer correction along the orbit of the pulsar, when it is located in a multiple system; $DM/f^2$ the delay for the signal propagation in the interstellar and interplanetary media at frequency $f$ taking into account the Doppler shift; and $\\Delta_{\\rm rel}$ the relativistic correction for the delay of the signal propagation in the gravitational field of the solar system.  The PATs at the solar-system barycenter are then used to calculate the rotational phase (number of rotations) of the pulsar: \\begin{equation}\\label{phase} N(T)=N_0+\\nu T +\\frac12\\dot\\nu{T}^2+\\varepsilon(T), \\end{equation} where  $N_0$ is the initial phase at time $T=0$, $\\nu$ and $\\dot\\nu$ are the rotational frequency of the pulsar and its derivataive at $T=0$, and $\\varepsilon(T)$ represents variations in the rotational phase (timing noise). The procedure for determining the parameters includes refining $\\alpha,\\delta,\\mu_\\alpha,\\mu_\\delta$ and $\\pi$, via a least-squares fit minimizing a weighted sum of the square differences between $N(T)$ and the nearest integer. The residual deviations in the rotational phase are usually expressed in units of the time $\\delta t=\\delta N/\\nu$. Here, we consider variations in the intrinsic rotation of the pulsar and variations due to irregularity of the reference time scale $\\Delta t_{clock}(T)$.  When intercomparing different realizations of atomic time \\cite{guinot88}, flicker noise in the frequency dominate on intervals of several months, while random walk noise in the frequency dominates on intervals of several years. Thus, clock variations have a power spectrum of the form $1/\\omega^{n}$ in the frequency domain, and are expressed in the time domain by the polynomial \\begin{equation} \\Delta {t}_{\\rm clock}({\\cal T})=c_0 + c{\\cal T}+\\frac 12\\dot c{\\cal T}^2+ \\frac 16\\ddot c{\\cal T}^3+\\ldots. \\end{equation} It is clear that the appearance of the quantity $\\Delta t_{\\rm clock}$ in Eq. (\\ref{phase}) leads to a redetermination of the rotational parameters of the pulsar: \\begin{equation} N({\\cal T})={N}_0'+(1+c)f {\\cal T}+\\frac 12(f\\dot c+(1+c)^2\\dot f){\\cal T}^2, \\end{equation} where $\\cal T$ is ideal barycentric time and $f$ and $\\dot f$ are the rotational frequency of the pulsar and its derivative undistorted by the clock parameters. For this reason, we must use the PAT values expressed in TT(BIPMXX), as best determined at the current time \\cite{guinot91}. ",
        "conclusions": "We have presented an algorithm for forming a group pulsar time based on optimal Wiener filtration. The algorithm makes it possible to distinguish the contributions to barycentric residual deviations of PATs due to irregularities in the intrinsic rotation of the pulsar and variations in the reference clocks.  Both irregularity in the pulsar rotation and variations in the reference time scale are obtained relative to an {\\it ideal} time system. Realization of the algorithm requires observations of at least two pulsars relative to a common reference time.  The proposed approach has better accuracy than the use of weighted averaging to form the group time scale, since it uses additional information about the signal via its correlation function or power spectrum.  The presence of a pulsar time that is independent of terrestrial conditions makes it possible to carry out independent checks of terrestrial time scales, which immediately provides possibilities for revealing the presence of variations that are common to all terrestrial standards that would otherwise be undetectable."
    },
    "1107/1107.3675_arXiv.txt": {
        "abstract": "The evolution of the interstellar medium (ISM) of elliptical galaxies experiencing feedback from accretion onto a central supermassive black hole has been studied recently with high-resolution 1D hydrodynamical simulations; these included cooling, heating and radiative pressure effects on the gas, specific for an average AGN spectral energy distribution, a RIAF-like radiative efficiency, mechanical energy, mass and momentum input from AGN winds, and the effects of starbursts associated with accretion.  Here we focus on the observational properties of the models in the soft and hard X-ray bands, specifically on 1) the nuclear X-ray luminosity; 2) the global X-ray luminosity and temperature of the hot ISM; 3) its temperature and X-ray brightness profiles, during quiescence, and before, during and after an outburst. After an evolution of $\\sim 10$ Gyr, the bolometric nuclear emission is very sub-Eddington ($l \\sim 10^{-4}$), and within the range observed, though larger than the most frequently observed values.  The nuclear bursts last for $\\approx 10^7$ yrs, and the duty-cycle of nuclear activity is a few$\\times (10^{-3}-10^{-2})$, when calculated over the last 6 Gyr.  The ISM thermal luminosity $\\lx$ oscillates in phase with the nuclear one, but drawing much broader peaks; a comparison with observed $\\lx$ values, for galaxies of optical luminosity similar to that of the models, shows that this behavior helps reproduce statistically the observed large $\\lx$ variation. At the present epoch, the largest observed $\\lx$ could be reproduced only by adding an external confining medium.  The average gas temperature is within the observed range; when limited to within $\\re$, its values lie on the upper half of those observed.  In quiescence, the temperature profile has a negative gradient; thanks to past outbursts, the brightness profile lacks the steep shape typical of inflowing models. After outbursts, disturbances are predicted in the temperature and brightness profiles (as analyzed also with the unsharp masking).  Most significantly, during major accretion episodes, a hot bubble from shocked hot gas is inflated at the galaxy center (within $\\approx 100$ pc); the bubble would be conical in shape, in real galaxies, and show radio emission. Its detection in the X-rays is within current capabilities, though it would likely remain unresolved.  The ISM resumes its smooth appearance on a time-scale of $\\approx 200$ Myr; the duty-cycle of perturbances in the ISM is of the order of 5-10\\%.  The present analysis reveals an agreement of the models with the observations, but also evidences that additional input physics is important in the ISM-black hole coevolution, to fully account for the properties of real galaxies. The main insertions to the models would be a confining external medium and a jet.  The jet will reduce further the mass available for accretion (and then $l$), and may help, together with an external pressure, to produce flat or positive temperature gradient profiles (that are common among galaxies in high density environments).  Alternatively to the jet, a reduction of $l$ can be obtained if the switch from high to low radiative efficiency takes place at a larger $l$ ($\\simeq 0.1$) than routinely assumed ($\\simeq 0.01$). ",
        "introduction": "\\label{intro}  Supermassive black holes (MBHs) at the centers of bulges and elliptical galaxies play an important role in the processes of galaxy formation and evolution (e.g., Cattaneo et al.~2009), as testified by remarkable correlations between host galaxy properties and the MBH masses (e.g., Magorrian et al.~1998, Ferrarese \\& Merritt 2000, Gebhardt et al.~2000, Graham et al.~2001) and as supported by many theoretical studies (e.g., Silk \\& Rees 1998, Haiman, Ciotti \\& Ostriker 2004; Merloni et al.~2004, Sazonov et al.~2005, Di Matteo, Springel \\& Hernquist 2005, Hopkins et al.~2006, Somerville et al.~2008, Kormendy et al.~2009).  An important aspect of the coevolution process is the radiative and mechanical feedback by the accreting MBH onto the galactic interstellar mediun (ISM) that is continuously replenished by normal stellar mass losses, at a rate of the order of $\\approx 1 M_{\\odot}$ yr$^{-1}$ in a medium-mass galaxy. In absence of feedback from a central MBH (and stripping from the intracluster medium in case of satellite galaxies), this ISM would develop a flow directed towards the galactic center, accreting $\\gsim 1 M_{\\odot}$ yr$^{-1}$ in a process similar to a ``cooling flow''.  Instead, in low mass galaxies, type Ia supernova (SNIa) heating is able to sustain a low-luminosity, global galactic wind (e.g., Ciotti et al.~1991, David et al.~1991, Pellegrini \\& Ciotti 1998), and the central MBH is in a state of permanent, highly sub-Eddington hot accretion (Ciotti \\& Ostriker 2011).  Therefore, in medium to high mass galaxies, feedback is required by the following empirical arguments: 1) the large amount of gas lost by the passively evolving stellar population during the galaxies' lifetime is not observed (e.g., Peterson \\& Fabian 2006), and just $\\lsim $1\\% of the mass made available by stars is contained in the masses of present epoch MBHs; 2) bright AGNs, as would be expected given the predicted mass accretion rate, are not commonly seen in the spheroids of the local Universe (e.g., Fabian \\& Canizares 1988, Pellegrini 2005). Thus AGN feedback is required just on the basis of a mass balance argument. Since over a large part of the galaxies extent, and for a large fraction of their lifetime, the ISM cooling time $t_{cool}$ is much lower than the galaxy age, one early quasar phase cannot be the solution to the cooling flow problem. The solution requires either steady heating, or heating with bursts on a timescale $\\Delta t\\approx t_{cool}$.  Unfortunately, on a purely theoretical ground, how much radiative and mechanical energy and momentum output from the MBH can effectively interact with the surrounding ISM, and what are the resultant MBH masses, is difficult to establish.  Recently, the interaction of the MBH output with the inflowing gas has been studied with high-resolution 1D hydrodynamical simulations in a series of papers (Ciotti \\& Ostriker 2007; Ciotti, Ostriker \\& Proga 2009, 2010 hereafter Papers I and III; Shin, Ostriker \\& Ciotti 2010a,b; Ostriker et al. 2010; Jiang et al.~2010), that are currently being extended to 2D treatments (Novak, Ostriker \\& Ciotti 2010).  These simulations implement a physically based detailed treatment of the radiative energy and momentum input from the MBH into the ISM, consistent both with observed average AGN spectra and theoretical calculations of radiation transport; they also include starformation, and a modelling of the mechanical energy and momentum feedback from AGN winds. The combined effects of radiative and mechanical feedback produce recurrent AGN burst phases accompanied by starformation, spaced apart by longer phases of relative quiescence. A cycle repeats with the galaxy seen alternately as an AGN/starburst for a small fraction of the time, and as a ``normal'' elliptical for much longer intervals. Accretion fueled feedback thus proves effective in suppressing long lasting cooling flows and in maintaining MBH masses within the range observed today, since the gas is mostly lost in outflows or consumed in starbursts.  Remarkably, {\\it while star formation is suppressed when the AGN is in the low-luminosity state, it is enhanced by the strong AGN outbursts, consistent with observations} (Schawinski et al. 2009).  Note finally that  a major role in producing global degassing, and in regulating the flow evolution, is also played by the SNIa's heating.  The previous papers, with the exception of Pellegrini, Ciotti \\& Ostriker (2009), were mainly dedicated to the study of the accretion physics and the feedback effects. Here instead we focus on the appearance that the models would have if observed in the X-ray band, both in quiescence and during an outburst of activity. We concentrate on two models (named \\bd and \\bt) extracted from the suite of cases presented in Paper III (Table 1 therein), characterized by a mechanical feedback efficiency dependent on the Eddington scaled accretion luminosity. \\bd and \\bt were considered particularly successful, since their input parameters agree with previous theoretical studies or observations (as, e.g., for the AGN wind opening angle, and the peak value of mechanical efficiency, that are in accord with those estimated from 2D and 3D numerical simulations; Proga, Stone \\& Kallman 2000, Proga \\& Kallman 2004, Benson \\& Babul 2009), and, at the same time, their final properties (as mass fraction of a younger stellar population, MBH mass, etc.) are in reasonable accord with observations.  The only difference in the input physics of \\bd and \\bt is the maximum value of the mechanical efficiency.  Since in the simulations the treatment of feedback is physically based, not tuned to reproduce observations, any agreement or discrepancy of the resulting model properties with X-ray observations is relevant to improving our understanding of the MBH-ISM coevolution, putting further constraints on the input ingredients, and possibly telling us what additional physics may be important in the problem. However, we stress that the models describe an isolated galaxy, where ram-pressure stripping (in case of satellite galaxies) and intracluster medium pressure confinement (in case of group or cluster central galaxies) are not taken into account (see Shin et al.~2010b).  The main observational signatures investigated here include the nuclear and gaseous emissions, and the ISM temperature and brightness profiles in the quiescent phases, and before, during, after a burst. Particular attention is paid to the appearance and detectability of central hot bubbles, with diameters of $\\sim$a hundred parsecs, that are produced by the models during outbursts, and to various kinds of disturbances in the hot ISM.  The analysis is performed in the soft and hard X-ray bands, also after unsharp-masking.  These predictions are relevant for their observational consequences, since the high angular resolution of the $Chandra$ satellite has allowed us to obtain the best definition ever for the hot gas properties of the galaxies of the local universe, by separating the contributions of stellar sources and hot gas, and the emission coming from different spatial regions within galaxies (e.g., Fabbiano 2011).  In particular, in several elliptical galaxies various kinds of hot gas disturbances have been detected, likely resultant from nuclear activity (e.g., Finoguenov \\& Jones 2001, Jones et al.~2002, Forman et al.~2005, Machacek et al.~2006, O'Sullivan et al.~2007, Million et al. 2010). At the same time, nuclear emission values have been detected down to very low luminosities, comparable to those of X-ray binaries (e.g., Loewenstein et al.~2001, Gallo et al. 2010, Pellegrini 2010).  The paper is organized as follows. Section 2 summarizes the main evolutionary phases of the representative models (\\bt and \\bt) considered; Section 3 describes how the observational properties of the models are derived; Section 4 presents a comparison of the nuclear luminosities with existing observations; Section 5 discusses the evolution of the X-ray luminosity and emission-weighted temperature of the ISM; Section 6 presents the projected temperature and surface brightness profiles at representative times during quiescence and a nuclear outburst. Finally, in Section 7 we summarize and discuss the main results. ",
        "conclusions": "The hot gas properties of massive ellipticals, with regard to luminosity and temperature, and their spatial distributions, allow us to derive insights into the hot gas evolutionary status, and its link with the host galaxy. Since the cooling times are short compared to the galaxy age, it is now commonly accepted that repeated cooling catastrophes have occurred in the past, accompanied by central starbursts and AGN outbursts.  The interest of a better understanding of this phenomenon is obvious, as it is directly linked to galaxy formation and evolution, and to the growth of the central MBH. Unfortunately, a complete theoretical picture is still missing, so that modeling coupled with a close comparison with observations is crucial in this field. In fact, these feedback events must leave signatures on the X-ray properties of the galaxies; indeed, the observed temperature and brightness profiles often cannot be fit easily with smooth profiles, or cooling flow profiles (e.g., Sarazin 2011, Statler 2011). In this work we have calculated the observational properties in the X-ray band of two galaxy models, representative of a class of detailed simulations of physically based models for the investigation of feedback modulated accretion in isolated galaxies. The feedback physics includes the combined effects of radiation and AGN winds (Paper III).  The observational properties derived provide good matches to what observed in general for the local universe, and account for a few otherwise puzzling observed properties; on the other hand, they also evidence the need for changes or additions to the input physics.  The main results of the present investigation are the following:  1) After an evolution of $\\simeq 10$ Gyr, the models are typically in a permanent quiescent phase.  The bolometric nuclear emission is very sub-Eddington ($l \\simeq 10^{-4}$), within the range observed for $l$, though the most frequently observed values ($l\\approx 10^{-5}$ or less) are somewhat lower. Unfortunately, uncertainties in the bolometric correction to apply to observed nuclear luminosities, appropriate for a spectral energy distribution at low emission levels, do not allow to make a stringent comparison between modelled and observed values.  However, also the nuclear X-ray emission $L_{\\rm BH,X}$ of the models, estimated as $L_{\\rm BH,X} \\lsim 0.2\\lbh$ as should be appropriate for low luminosity nuclei, and $L_{\\rm BH,X}/\\ledd$, tend to lie on the upper end of what is observed.  Thus in real galaxies an additional mechanism may be required to reduce further the mass available for accretion; this could be provided by the mechanical feedback of a (nuclear) jet, and/or by a thermally driven wind from a RIAF. Alternatively, the switch from disc to RIAF behavior (that is $l\\propto \\dot m^2$) should occur at a larger $l$ than assumed here (e.g., at $l\\simeq 0.1$ rather than $l\\simeq 0.01$). Simulations of ram-pressure stripping effects, instead, showed that a reduction in the accretion rate is not attained, because in the quiescent hot accretion regime, the accreting mass comes mainly from the stellar mass losses within the central $\\sim 100$ parsecs of the galaxy, that are not affected by stripping (Shin et al. 2010b).  2) The X-ray luminosity of the ISM oscillates in phase with the nuclear luminosity, though with much broader peaks; at the present epoch, $\\lx$ lies at the lower end of the large observed range for galaxies of $\\lb$ similar to that of the models.  The degassing/heating in the models may then be too efficient, or a larger/more concentrated gravitating mass, or a confining external medium, are needed.  However, when the gas luminosities during the whole evolution are considered, the observed $\\lx$ range is better reproduced.  This is even more true if also an additional model of larger $\\sigma$, and same $\\lb$, is included.  Part of the observed large variation in $\\lx$ for galaxies of a given $\\lb$ could then be explained by nuclear activity.  3) The average ISM temperature is within the observed large range for the model $\\sigma$; when estimated within $\\re$, the model temperatures reproduce better the upper half of those observed. Modifications to the models by the addition of a jet or an external medium, as suggested in the previous points, should then not increase the average temperature within $\\re$.  4) During quiescence, the profiles of the gas temperature and brightness resemble those observed for many local galaxies. Especially remarkable is the lack of the steep brightness profile shape typical of inflowing models, due to the frequent removal of gas from the galactic central regions, and to the heating provided by mechanical feedback (that is always present, even during quiescent phases).  The models show negative temperature gradients, that are common for isolated galaxies; the addition of a jet or a confining agent should change the temperature profiles into a flat or outwardly increasing profile, as also frequently observed for galaxies in rich environments.  5) During outbursts, disturbances are predicted in the temperature and brightness profiles; the ISM resumes the smooth appearance of steady and low-luminosity hot accretion on the MBH on a time-scale of $\\lsim 200$ Myrs.  The most conspicuous variations with respect to smooth profiles are within current detection capabilities, and could correspond to (part of) the widespread disturbances observed in galaxies of the local Universe. In particular, shocked hot gas should be seen at the galactic center (within $\\sim 100$ pc), possibly not resolved; this would be a certain sign of prior AGN activity.  These hot bubbles could be revealed by emission of cosmic-rays, in a structure similar to a gigantic supernova remnant. Preliminary results from 2D runs show bipolar nuclear outflows, that should be seen as conical cavities extending from the galactic center, and may be called jet-like features.  6) The duty-cycle of nuclear activity is of the order of a few $\\times (10^{-3}-10^{-2})$, depending on the assumed mechanical feedback efficiency; in general, a burst cycle lasts for $\\approx 10^7$ yrs. These duty-cycle values are broadly consistent with the fraction of active galaxies measured in observational works, though reported values for the local universe are somewhat lower, for the MBH mass of the models.  In order to make a more consistent comparison with observations, the dataset of models should be increased, and the duty cycle computed only for the last 2--3 Gyrs; this will reduce the duty cycle, as the models are characterized by a declining nuclear activity.  7) The duty-cycle of perturbances in the ISM is of the order of 5-10\\%, from their average number and duration.  This duty-cycle likely increases with galaxy mass, because an outburst has a greater impact in less massive (and less gas-rich) systems, which then are \"on\" for a shorter time (Ciotti \\& Ostriker 2011). The ISM duty-cycle, and its trend with galaxy mass, compare reasonably well with preliminary estimates obtained from a large sample of hot gas coronae in elliptical galaxies observed with $Chandra$ (Nulsen et al. 2009): the fraction of galaxies with X-ray cavities in the hot gas is $\\lsim 10$\\% when $\\lx<10^{41}$ erg s$^{-1}$ (as for the models), and reaches $\\sim 25$\\% in the most luminous ones. The presence of cavities has been attributed to the action of jets inflating radio lobes and displacing the surrounding gas. Cavities can also be created in the scenario presented here, as hinted for by 2D simulations, due to bipolar nuclear outflows.  7) Two diagnostic planes have been constructed. In the first one, the nuclear luminosity $\\lbh$ and the ISM luminosity $\\lx$ are followed during the whole model evolution. The points representative of the models populate a wedge region, that should then be occupied when observing a large set of galaxies.  Another plane shows the evolution of $\\lx$ versus the average gas temperature $\\ta$; here the most populated region is that of a large $\\lx$ variation (factor of $\\sim 10$) for $\\ta$ keeping between 0.4 and 0.6 keV.  Clearly, a larger set of models is to be explored, in order to better establish the final gas properties (as gas content, nuclear and ISM duty-cycle, etc.) to be compared with those of a statistically large sample.  For example, a general expectation is that changes of the galaxy properties have an impact on the number of nuclear outbursts: depending on many model parameters (supernova rate, central $\\sigma $, dark matter amount and distribution, and even external pressure due to an intragroup or intracluster medium), bursts could take place even towards the present epoch, or be confined to the early epoch (Ciotti \\& Ostriker 2011)."
    },
    "1107/1107.1994_arXiv.txt": {
        "abstract": "{ We report on a sensitive search for H$_2$ 1-0 S(1), 1-0 S(0) and 2-1 S(1) ro-vibrational emission at 2.12, 2.22 and 2.25 $\\mu$m in a sample of 15 Herbig Ae/Be stars employing  CRIRES, the ESO-VLT near-infrared high-resolution spectrograph, at R$\\sim$90,000. We report the detection of the H$_2$ 1-0 S(1) line toward HD 100546 and HD 97048. In the other 13 targets, the line is not detected. The H$_2$ 1-0 S(0) and 2-1 S(1) lines are undetected in all sources. These observations are the first detection of near-IR H$_2$ emission in HD 100546. The H$_2$ 1-0 S(1) lines observed in HD~100546 and HD~97048 are observed at a velocity consistent with the rest velocity of both stars, suggesting that they are produced in the circumstellar disk. In HD~97048, the emission is spatially resolved and it is observed to extend at least up to 200 AU from the star. We report an increase of one order of magnitude in the H$_2$ 1-0 S(1) line flux with respect to previous measurements taken in 2003 for this star, which suggests line variability. In HD 100546 the emission is tentatively spatially resolved and may extend at least up to 50 AU from the star. Modeling of the H$_2$ 1-0 S(1) line profiles and their spatial extent with flat keplerian disks shows that most of the emission is produced at a radius larger than 5 AU. Upper limits to the H$_2$ 1-0 S(0)/ 1-0 S(1)  and H$_2$ 2-1 S(1)/1-0 S(1) line ratios in HD 97048 are consistent with H$_2$ gas at T$>$2000 K and suggest that the emission observed may be produced by X-ray excitation. The upper limits for the line ratios for HD~100546 are inconclusive. Because the H$_2$ emission is located at large radii, for both sources a thermal emission scenario (i.e., gas heated by collisions with dust) is implausible. We argue that the observation of H$_2$ emission at large radii may be indicative of an extended disk atmosphere at radii $>$ 5 AU. This may be explained by a hydrostatic disk in which gas and dust are thermally decoupled or by a disk wind caused by photoevaporation. } ",
        "introduction": "Circumstellar disks around pre-main sequence stars are the birth places of planets. Hence, the characterization of their physical properties is of paramount importance. At the time when giant planets are in formation, protoplanetary disks are rich in gas. However, there are relatively few observational constraints of the gas content in the disk, in particular for the inner disk ($R<10$ AU), the region where planets are expected to form. To study the gas content in the disk,  molecular and atomic line emission is used. Because the disk has a radial temperature gradient, different transitions of diverse gas tracers probe different radii in the disk (see reviews by Najita et al. 2007 and Carmona 2010). These emission lines are in general produced in the surface layers of the disk where the dust is optically thin.  By far the main constituent of the gas in protoplanetary disks is molecular hydrogen (H$_2$). However, given its physical nature (H$_2$ is an homonuclear molecule that lacks a dipole moment), H$_2$ transitions are very weak. In consequence,  in contrast to other gas tracers (most notably CO), H$_2$ emission is harder to detect. With the advent of space borne observatories in the UV and infrared and ground based high-resolution infrared spectrographs, it has been possible to start the search and study of H$_2$ emission from disks around young stars from the  UV  to the mid-IR.  In the UV, H$_2$ electronic transitions trace in emission hot H$_2$ gas in the disk that is photoexcited (``pumped\") by Ly$\\alpha$ photons (e.g. Valenti et al. 1993; Ardila et al. 2002; Herczeg et al. 2006) or excited by electrons that are generated by X-rays (e.g., Ingleby et al. 2009, France et al. 2011). In absorption, UV transitions trace cold, warm, and hot H$_2$ in the line of sight (e.g., Martin-Za\\\"idi et al. 2005, 2008a), either from the disk itself (e.g., flared disk, close to edge-on disk, disk wind) or the star's circumstellar environment (e.g., envelopes). In the mid-IR,  H$_2$ pure-rotational transitions trace thermal emission of warm gas at  few hundred K (e.g. Thi et al. 2001; Bitner et al. 2007; Martin-Za\\\"idi et al. 2007, 2008b, 2010; Lahuis et al. 2007; Carmona et al. 2008b). In the specific case of the near-infrared, ro-vibrational transitions of H$_2$ trace thermal emission of hot H$_2$ at a  thousand K, or H$_2$ gas excited by UV or X-rays (e.g. Weintraub et al. 2000; Bary et al. 2003, 2008; Itoh et al. 2003; Ramsay Howat \\& Greaves 2007; Carmona et al. 2007, 2008a). Because gas at these temperatures is expected to be located at radii up to a few AU, H$_2$ near-IR emission has the potential of tracing the gas in the terrestrial planet region of disks if the observed emission is thermal.  H$_2$ near-IR emission from disks has been mainly studied toward the low-mass T Tauri stars, where it has been detected preferentially in objects exhibiting signatures of a high accretion rate (i.e., large H$\\alpha$ equivalent widths and $U-V$ excess) and in a few weak-line T Tauri stars with bright X-ray emission (see the discussion sections of  Carmona et al. 2007, 2008a, and Bary et al. 2008). However, for intermediate mass young-stars, i.e., Herbig Ae/Be stars, there is to date only one reported detection of near-infrared H$_2$ emission at the velocity of the star, namely in the star HD 97048 (Bary et al. 2008). Because Herbig Ae/Be stars are bright in the near-IR, the detection of the faint H$_2$ lines is challenging, in particular at low spectral resolution. In this paper, we present the results of a considerable effort that we undertook to search for H$_2$ 1-0 S(1), 1-0 S(0) and 2-1 S(1) emission toward Herbig Ae/Be stars employing  CRIRES, the ESO-VLT near-infrared high-resolution spectrograph, at resolution $R\\sim$ 90,000.\\\\  The paper is organized as follows. In Section 2 we describe our CRIRES observations and the data reduction techniques. In Section 3 we present the resulting spectra and constrain the excitation mechanism of the detected lines based on H$_2$ line ratios. We then analyze the observed H$_2$ lines in the context of a flat disk model. In Section 4 we discuss our results. First, we compare the H$_2$ line profiles detected with the line profiles of other disk gas tracers. We then discuss our H$_2$ observations and observations of dust in the inner disk. Thereafter, we analyze H$_2$ near-IR emission in the frame of diverse disk structure models. We propose explanations for the relatively large number of non-detections, and explore diverse scenarios to explain the detections of near-IR H$_2$ lines in HD 100546 and HD 97048. Finally, in Section 5 we present our conclusions. ",
        "conclusions": "  {\\it a)} The non-detections of near-IR H$_2$ can be explained {\\it (i)} as a natural consequence of a passive disk structure in which the gas and dust are thermally coupled (i.e., the amount of optically thin hot H$_2$ gas in the inner disk is too low to produce detectable levels of H$_2$ emission). Or  {\\it (ii)}  in the context of a passive disk with thermally decoupled dust and gas; here the H$_2$ close to the star is photo-dissociated, and temperatures in the outer disks of self-shadowed disks are too low to produce H$_2$ emission.  {\\it b)} Most of the H$_2$ emission observed in HD~100546 and HD~97048 is produced at $R>>5$ AU and extends at least to $R>50-200$ AU. This location is much farther out in the disk than the radii up to a few AU where temperatures in the disk are expected to reach the 1000-3000 K that is needed to excite the observed H$_2$ lines thermally in standard passive disk models with thermally coupled gas and dust. The emission can be explained by {\\it (i)} disk models with thermally decoupled gas and dust that allow for high gas temperatures at large radii in flaring disks. At small radii, H$_2$ is photo-dissociated by the strong stellar UV field and H$_2$ formation efficiencies are low. After $\\sim$ 10 AU, H$_2$ formation becomes more efficient at balancing  photo dissociation, thus leading to detectable H$_2$ gas in the disk atmosphere; or {\\it (ii)} by  an extended disk atmosphere owing to a photoevaporative disk wind. More detailed modeling would be of great help to understand the origin of the H$_2$ near-IR emission observed  in HD~100546 and HD~97048."
    },
    "1107/1107.4864_arXiv.txt": {
        "abstract": "{Stellar activity induces apparent radial velocity (RV) variations in late-type main-sequence stars that may hamper the detection of low-mass  planets and the measurement of their mass.} {We use simultaneous measurements of the active planet host star HD~189733 with high-precision optical photometry by the MOST satellite and high-resolution spectra by SOPHIE. We apply on this unique dataset a spot model to predict the activity-induced RV variations and compare them with the observed ones.} {The model is based on the rotational modulation of the stellar flux. A maximum entropy regularization is applied to find a unique and stable solution for the distribution of the active regions versus stellar longitude. The RV variations are synthesized considering the effects on the line profiles of the brightness perturbations due to dark spots and bright faculae and the reduction of the convective blueshifts in the active regions. } {The synthesized RV time series shows a remarkably good agreement with the observed one although variations on timescales shorter than $4-5$ days cannot be reproduced by our model. Persistent active longitudes are revealed by the spot modelling. They rotate with slightly different periods yielding a minimum relative amplitude of the differential rotation of $\\Delta \\Omega / \\Omega = 0.23 \\pm 0.10 $. Moreover, several active regions with an evolution timescale of $2-5$ days and an area of $0.1-0.3$ percent of the stellar disc are detected. } {The method proves capable of reducing the power of the activity-induced RV variations by a factor from 2 to 10 at the rotation frequency and its harmonics up to the third. Thanks to the high-precision space-borne photometry delivered by CoRoT, Kepler, or later PLATO, it is possible to map the longitudinal distribution of active regions in late-type stars and apply the method presented in this paper to reduce remarkably the impact of stellar activity on their RV jitter allowing us to confirm the detection of low-mass planets or refine the measurement of their mass.} ",
        "introduction": "Several studies have recently addressed the perturbation of the RV induced by stellar activity.  This jitter may severely hamper the detection of low-mass planets or the measurement of the mass of those found by the method of the transits.  \\citet{Lagrangeetal10}, \\citet{Meunieretal10a}, and \\citet{Meunieretal10b} have considered the Sun as a star finding that the reduction of convective blueshifts in active regions, due to the quenching of turbulent motions by the   photospheric magnetic fields, has a remarkable effect inducing variations with an amplitude up to $5-10$~m~s$^{-1}$. \\citet{Dumusqueetal11} have extended those studies introducing a model for the impact of dark spots on the RV variations of solar-like stars based on the statistics of sunspot groups.  Several techniques have been proposed to reduce the impact of the RV perturbation induced by stellar activity. They can be classified into two broad groups. Those in the first group exploit the correlation between the RV perturbation and other properties of the cross-correlation function used to measure the RV itself, such as the mean line bisector shift and/or its FWHM \\citep[e.g., ][]{Meloetal07,Boisseetal09,Pontetal11}. The main limitation of such techniques is the lack of a tight correlation between the RV perturbation and those proxies for the average spectral line distortion, as demonstrated by, e.g., \\citet{Boisseetal11}. The other group consists of the techniques that exploit the modulation of the activity-induced perturbation by the stellar rotation. Most of the power due to the activity-induced perturbation is indeed concentrated at the rotation frequencies and its first two or three harmonics. Therefore,  it is possible to filter it out leaving the signal due to the reflex orbital motion induced by a planetary companion, unless its orbital period falls close to such frequencies, and to fit simultaneously activity and planetary signals \\citep[see ][ for details]{Boisseetal11}. This approach has been recently applied, among others, to detect a second non-transiting planet around the active star \\object{CoRoT-7} and to refine the measurement of the mass of its transiting telluric planet CoRoT-7b \\citep[see ][]{Quelozetal09,Boisseetal11}.  In addition to those methods, another approach to reduce the impact of stellar activity can be based on the models  of the active region distribution derived from  high-precision photometry such as that provided by the space-borne telescopes of MOST, CoRoT, or Kepler for their target stars. From their light curves, it is possible to derive the longitudes of the active regions and their area, tracing their evolution along successive rotations. This information can be used to model  the RV perturbation induced by stellar activity. \\citet{Lanzaetal10} have applied this approach to the  light curve of CoRoT-7 to predict the power spectrum of the RV perturbation induced by stellar activity and derive the significance of the signals attributed to the telluric planets CoRoT-7b and CoRoT-7c. The main limitation of that study was the lack of simultaneity between the CoRoT light curve and the HARPS RV time series analysed to extract the reflex motions induced by the two planets.  Here we further explore this approach by means of simultaneous photometric and RV measurements of HD~189733 making use of the MOST light curve and the SOPHIE RV time series published by \\citet{Boisseetal09}. They span about one month in the summer of 2007, corresponding to two full rotations of the star with a period of $P_{\\rm rot} \\sim 12$~days \\citep{HenryWinn08}. \\object{HD~189733} is accompanied by a Jupiter-mass transiting planet with an orbital period of 2.218 days. With an effective temperature of $\\sim 5000$~K { and an average chromospheric index $S_{\\rm HK} = 0.525$}, this K1 main-sequence star is remarkably active \\citep{Bouchyetal05,Moutouetal07,Boisseetal09}. Its radial velocity (hereafter RV) measurements are then perturbed with variations up to $\\sim 20-30$~m~s$^{-1}$, as evident in the residuals of the RV  best fit of the planetary orbit \\citep{Boisseetal09}.  Our spot modelling technique allows us to retrieve the location and the evolution of the photospheric active regions of HD~189733 finding also evidence of a remarkable stellar differential rotation, as recently found by \\citet{Faresetal10} by means of Doppler Imaging techniques. The information on the starspot longitude and evolution is used to synthesize the RV variations that can be compared with the simultaneous SOPHIE measurements to assess the advantages and limitations of the technique. We find that the modulation of the RV perturbation induced by stellar activity is reproduced remarkably well on timescales comparable with the stellar rotation period but it is not possible to model the variations with a timescale of $1-2$ days, i.e., much shorter than the stellar rotation period. This is an unavoidable consequence of the fact that the rotational modulation of the optical flux is used to map starspots. On the other hand, the main advantage of this method is the remarkable reduction of the perturbations at frequencies comparable with the stellar rotation frequency and its low-order harmonics, that is an important improvement over previously proposed approaches.   The next section describes the observational dataset. Then, we define our model: 1) how active regions locations are derived from lightcurve; 2) how we compute the distortion of the line profiles due to the brightness perturbations owing to dark spots and bright faculae and the reduction of the convective blueshifts in the active regions; and 3) how the distributions of active regions derived from the spot model fitted to the light curve are used to synthesize RV variations. The fourth section gives the parameters chosen for our modelling of HD~189733. The results, given in  Section 5, are discussed in the final section. ",
        "conclusions": "\\label{conclusions}   We have applied the ME spot modelling method introduced by \\citet{Lanzaetal09a} and \\citet{Lanzaetal09b} to the MOST photometry of  HD~189733. It is among the brightest and most active stars accompanied by a transiting hot Jupiter and displays activity-induced RV variations  with an amplitude reaching  several tens of m~s$^{-1}$. The simultaneous RV measurements collected with SOPHIE offer us a unique opportunity to compare the models proposed to account for the RV variations induced by activity with the observations. % Both the MOST photometry and the SOPHIE RV measurements have errors comparable with those expected with larger telescopes on fainter stars. In the case of CoRoT or Kepler, it is possible to reach a photometric relative accuracy of $1.5 \\times 10^{-4}$ and $2.0 \\times 10^{-5}$ on a G2V star of magnitude $V=12$ in one hour integration time, respectively. Using HARPS, we can improve the radial velocity accuracy down to $2-5$ m~s$^{-1}$ on stars of that apparent magnitude, making our test particularly relevant for those targets, especially for the most active ones, given that HD~189733 is among the most active planetary hosts found so far.  Our spot models show several active regions with lifetimes comparable or longer than  the duration of  MOST observations, i.e., $\\sim 30$ days, that evolve remarkably and migrate at different rates with respect to a reference frame rotating with the mean stellar rotation period. This is interpreted as evidence of surface differential rotation with a relative amplitude of $\\Delta \\Omega / \\Omega = 0.23 \\pm 0.10$, in agreement with the result obtained from Doppler Imaging \\citep{Moutouetal07,Faresetal10}. It is interesting to note that, although the  MOST observations cover only $\\sim 2.5$ stellar rotations, the differential rotation signal is fairly evident thanks  to its large amplitude. Moreover, a population of small active regions with a surface of $0.1-0.3$ percent of the stellar disc and typical lifetimes of $1-2$ days is observed, as indicated by the  oscillations in the residuals of the photometric model  (cf. Fig.~\\ref{long_distr}).  We introduce a method to synthesize the RV variations induced by stellar activity from the longitude distribution of the stellar active regions derived by the ME spot modelling. The model assumes that the stellar active regions have the same contrasts and turbulent convection properties of the solar active regions. Starspots are assumed to be close to the stellar equator because we cannot derive  their latitudes from the high-precision photometry. Indeed, in the case of HD~189733  the spot latitudes could be estimated from the migration of the  active longitudes, given the differential rotation law of \\citet{Faresetal10}, but the results are uncertain and we prefer to compute the RV variations adopting nearly equatorial spots to test our approach in the general case when the latitudinal dependence of the rotation period is unknown.  Our model has only two free parameters, i.e., an offset  applied to match the zero points of the synthesized and observed RV scales, and a coefficient $\\langle C_{\\rm d} \\rangle $ measuring the average reduction of the depth of the spectral lines in dark spots. They are derived by a best fit to the RV observations. The synthesized RV variations obtained from the starspot distributions  show a remarkably good agreement with the simultaneously observed RV variations, although the latter display remarkable changes on timescales as short as $1-2$ days that cannot be accounted for by our model whose spot distributions can be updated only with a time interval of $\\sim 4$ days. This is a limitation of every model based on the rotational modulation of spot visibility because it requires a minimum time interval, not too short in comparison with the stellar rotation period, to map the brightness inhomogeneities on the stellar surface.  The RV variations set also a constraint on the facular effects suggesting that dark spots are dominating both the photometric and the activity-induced RV variations, at least during the time interval covered by the present observations. Of  course,  this does not exclude the possible presence of bright faculae in the photosphere of HD~189733, but their contribution cannot be detected in the available data. We show that the best value of $Q$, the facular-to-spotted area ratio, is $Q=0$, with an acceptable range extending from~$\\sim 0$~to~$\\sim 5$ when only the modelling of the optical flux variations is considered. Considering the modelling of the RV variations, an upper limit of $Q \\sim 2-3$ is derived.  In the case of the Sun, the best value of $Q$ is $9$ \\citep{Lanzaetal07}. Thus our results indicate a lower relative contribution of the faculae to the  light variation  of HD~189733. The amplitude of the rotational modulation of our star was $\\sim 0.015$ mag during the present MOST observations and $\\sim 0.03$ mag during the 2006 observations described in \\citet{Miller-Riccietal08}, i.e., $\\sim 5-10$ times that of the Sun at the maximum of the eleven-year cycle. This indicates that HD~189733 is remarkably more active than the Sun, which may account for the undetectable facular contribution to its light variations, as suggested by \\citet{Radicketal98} and \\citet{Lockwoodetal07}.  We conclude that the approach presented in this paper can reduce the power of the activity-induced RV variations at the rotation frequency and its low-order harmonics by a factor ranging from 2 to 10, which represents an important improvement for the detection of exoplanets corotating with the star or with an orbital frequency in the $2:1$ ratio  with the rotation frequency, as suggested by the statistical studies of \\citet{Lanza10} and \\citet{Lanzaetal11}. Thanks to the high-precision space-borne photometry already available with MOST, CoRoT, and Kepler, it is possible to map the longitudinal distribution of active regions in late-type stars and apply the method presented in this paper to reduce remarkably the impact of stellar activity on their RV jitter  allowing us to confirm the detection of telluric planets or refine the measurement of their mass \\citep[see, e.g., the application to CoRoT-7 in ][]{Lanzaetal10}."
    },
    "1107/1107.4235_arXiv.txt": {
        "abstract": "{ During the past few years, first observations of starburst galaxies at $>$~GeV energies could be made with the Fermi Gamma-ray Space Telescope (GeV range) and Imaging Air Cherenkov Telescopes (TeV range). The two nearest starbursts, M82 and NGC253 were detected, and most recently, the detection of two starburst-Seyfert composites (NGC1068 and NGC4945) were reported. The emission for the two starbursts is best explained by hadronic interactions, and thus providing a first, unique opportunity to study the role of cosmic rays in galaxies. In this paper, the role of cosmic rays for the non-thermal component of galaxies is reviewed by discussing the entire non-thermal frequency range from radio emission to TeV energies. In particular, the interpretation of radio emission arising from electron synchrotron radiation is predicted to be correlated to TeV emission coming from interactions of accelerated hadrons. This is observed for the few objects known at TeV energies, but the correlation needs to be established with significantly higher statistics. An outlook of the possibility of tracing cosmic rays with molecular ions is given. ",
        "introduction": "Galaxies with a high star formation rate provide the opportunity to study the influence of massive stars on the large-scale behavior of galaxies in detail. In particular, the role of gas heating by star formation can be studied by the observation of far-infrared emission and the role of supernova remnants can be investigated by looking at non-thermal radio emission from electron synchrotron radiation. Most recently, first gamma-ray detections using the Fermi Gamma-ray Space Telescope (FGST) and Imaging Air Cherenkov Telescopes like H.E.S.S.\\ and VERITAS were made. The nearest starburst galaxies M82 and NGC253 were detected above GeV energies, tracing the interaction of hadronic cosmic rays with the gas in the galaxy.  In this paper, the multifrequency spectrum of starburst galaxies will be reviewed. In Section \\ref{radio_fir}, the correlation between the far infrared and radio emission is discussed. In Section \\ref{proton_proton:sec}, hadronic interactions at energies $E>$~GeV are discussed in the context of neutral secondary production. In Section \\ref{crtracers:sec}, cosmic ray-induced ionization is mentioned as a future method of tracing hadronic cosmic rays in galactic environments. In particular,  a supernova remnant interacting with a molecular cloud provides an optimal environment for the production of H$_{2}^{+}$, at a level which should be detectable in the future with instruments like Herschel and ALMA. ",
        "conclusions": ""
    },
    "1107/1107.3881_arXiv.txt": {
        "abstract": "{In this paper, a grid of the binary evolution models are calculated for the study of blue straggler (BS) population in intermediate age ($\\log Age$=7.85-8.95) star clusters. The BS formation via mass transfer and merging is studied systematically using our models. Both Case A and B close binary evolutionary tracks are calculated in a large range of parameters. The results show that BSs formed via Case B are generally bluer and even more luminous than those produced by Case A. Furthermore, the larger range in orbital separations of Case B models provide a probability of producing more BSs than Case A. Based on the grid of models, several Monte-Carlo simulations of BS populations in the clusters in the age range are carried out. The results show that BSs formed via different channels populate different areas in color magnitude diagram(CMD). The locations of BSs in CMD for a number of clusters are compared to our simulations as well. In order to investigate the influence of mass transfer efficiency in the models and simulations, a set of models are also calculated by implementing a constant mass transfer efficiency, $\\beta$=0.5 during Roche lobe overflow (Case A binary evolution excluded). The result shows BSs can be formed via mass transfer at any given age in both cases. However, the distributions of the BS populations on CMD are different. ",
        "introduction": "\\label{sect:intro}  Blue stragglers are observed in almost all stellar systems which appear to be anomalously young compared to other populations in the system. They are commonly defined as stars brighter and bluer than the turn-off(TO) point on CMD. Such stars seem to remain on the main sequence(MS) longer than is estimated by standard theory of stellar evolution. The existence of BSs can significantly affect the integrated spectral energy distributions, especially  in blue and UV bands (Li \\& Han 2009, Xin, Deng \\& Han 2007), thus challenges traditional single stellar population(SSP).  Ferraro et al.(1993) found two groups of BSs in M3, which have a bimodal radial distribution. They concluded whether the two groups of BSs have different physical properties or they represent the products of different formation mechanisms. Laget et al.(1994) identified 24 BSs in a UV imaging study of M3. They found two distinct classes of BSs in M3. One has normal UV flux distribution and may be the massive stars formed through stellar collisions, whereas the other has UV excess and appears to be composite systems consisting of a hot secondary star orbiting a MS turn-off star or a subgiant or even possibly a low luminosity giant. Recently, Ferraro et al.(2009) observed two distinct sequences of BS populations on the CMD of M30. They concluded that the bluer sequence is arising from direct collisions and the redder one arising from the evolution of close binaries possibly still experiencing mass exchange. All the observations above suggest that BSs may have different formation mechanisms.  Though BS formation mechanisms are not completely understood yet, many possible scenarios now have been proposed to explain BS formation (see the review of Stryker 1993). At present, some possible ways are believed to account for the BS formation: mass transfer (MT) or complete coalescence in close binaries; stellar mergers resulting from direct collisions. It is believed that more than one scenarios are necessary to explain the observations. However, the relative importance of the two possible mechanisms remains unclear.  The mechanism of mass transfer in close binary was first proposed by McCrea (1964). Mass transfer can occur when the massive component in the binary system expands out of the Roche lobe. As a result of mass transfer, the companion can become a more massive main-sequence star than turn-off stars whose life time could be notably extended, or even doubled compared to a star with the same final mass. Coalescence of the two components may also occur during the process. Three mass transfer cases (case A, case B and case C) are defined according to the evolutionary state of the primary at the onset of mass transfer. They are associated with the evolutionary stages respectively: case A on the main sequence, case B after the main sequence before helium ignition and case C during central He-burning and afterwards (Kippenhahn \\& Weigert 1967). The mass of BS formed this way cannot exceed twice the cluster turn-off mass. Besides blue stragglers, many other observational objects are also associated with interactive binaries, such as Ba stars which can be explained by mass transfer through wind accretion, disk accretion or common envelope ejection (Han et al. 1995; Liu et al. 2009). Binary interactions are sometimes linked to variable stars observed in star clusters. It is known that some cluster variable stars are still in binary systems, such as symbiotic stars, W UMa stars and EA-type binaries. A series of work searching for variables in clusters have been performed for now (Lu et al. 2006,2007; Xin et al. 2002; Zhang et al. 2003,2005; Li et al. 2004). Moreover, binaries (white dwarf + MS or red giant) are believed to be possible progenitors of type Ia supernova (Meng et al. 2006,2009; Wang et al. 2008,2010; Guo et al. 2008).  BSs could also, in principle, be formed in direct collisions between main sequence stars. Direct collision hypothesis was originally presented by Hills \\& Day (1976). They proposed that the remnant of two main-sequence stars collision could also produce a blue straggler. It is likely the cause of making BSs in globular clusters, where star densities are extremely high and dynamical effects are thought to play a dominant role. However, it is unlikely to happen for open clusters, as the time-scale of collision is too long (Press \\& Teukolsky 1977; Mardling \\& Aarseth 2001) and neither the peculiar compositions of Ap and Bp-type blue stragglers, nor different rotation rates of observed BSs are by themselves sufficient to explain the large color differences between observed BSs and MS stars (Mermilliod 1982).  Different mechanisms are believed to dominant in different environments. As a consequence of dynamic evolution of the host cluster, the collision scenario is thought to be responsible for BSs in dense cluster cores. Binary interaction process and products are believed to dominate in more spare environment such as open clusters and in the field (Mapelli et al. 2004; Lanzoni et al. 2007; Dalessandro et al. 2008; Sollima et al. 2008).  Many detailed work about the two mechanisms is presented. Pols \\& Marinus (1994) performed Monte-Carlo simulations of close binary evolution in young clusters. The BSs predicted by their work agree well with the observations in quantitatively for clusters younger than about 300 Myr but are not enough for the clusters of older ages (between 300 and 1500 Myr). N-body simulations of BSs in old open cluster M67 (Hurly et al. 2001,2005) suggest that the formation of BSs is dominated by both mass transfer and cluster dynamics.  Recently, the binary explanation was claimed to be dominant even in dense globular cluster cores (Knigge \\& Sills 2009). They found a clear, but sub-linear correlation between the number of BSs in a cluster core and the total stellar mass contained within it. It is regarded as the strongest and most direct evidence to date that most blue stragglers, even those found in cluster cores, are the progeny of binary systems.  Some other mechanisms can also contribute to BS population. BSs predicted in mass transfer theory are fainter than 2.5 mag above the cluster turn-off stars, but there are exceptions in observations such as F81 in M67. To explain these massive stragglers, Hagai and Daniel (2009) discussed the possibility for BS formation in primordial and/or dynamical hierarchical triple stars. Chen \\& Han (2008) estimated the remnants of binary coalescence from case A models. The outcome BSs in their work can also be able to predict some high luminous BSs, however, the small quantity can hardly match the luminous BSs in M67. Angular momentum loss (AML) in low mass binaries can also contribute to BS populations, but only in old clusters. There are a number of subjects including the treatment of AML (Li et al. 2004; Demircan et al. 2006; Micheal \\& Kevin 2006; Stepien 2006). The recent study of Chen \\& Han (2008a,2009) also demonstrated that AML is likely a main factor to BS formation in old open clusters. However, its contribution can be ignored in clusters younger than 1.0 Gyr.  It is more difficult to study BSs in young clusters than in old ones in a systematic manner, as the observed BSs in younger clusters are usually very limited due to the relative small number of massive main sequence stars. We investigated case A mass transfer in close binaries in young clusters NGC 1831, whose log(age) is around 8.65 by isochrone fitting. The result shows that binaries experiencing case A mass transfer can hardly account for the luminous BSs in such young clusters (Lu \\& Deng 2008). A subsequent work combined case A and B is then performed to study BSs in old open clusters. The evolutionary behaviors of case A and B are compared in detail. We concluded that case B is as important as case A in old clusters, and it is even more important in younger clusters by comparing the different evolutionary behaviors with different initial primary mass. We predicted that BSs formed via case B mass transfer can be brighter and bluer than case A, and the number of BSs can be remarkably increased (Lu et al. 2010, hereafter LDZ10) thus provide a clue to study the statistics of BSs in such young clusters.  In this work, we apply the same approach as in LDZ10 to systematically study BS populations in younger stellar populations represented by open clusters. We select clusters younger than 1.0 Gyr in this paper, because stellar densities in young clusters are sufficiently low that dynamical effects are not so important to complicate our study. Furthermore, there is good evidence for the presence of a significant fraction of primordial binaries in such young clusters, making it not only interesting but also essential to better understand the evolution of the cluster. A large grid of binary evolutionary calculations with the initial mass of the donor from 2.0-7.6 $M_\\odot$ are carried out in this work. The models can cover all close binaries experiencing mass transfer between 70.0 Myr and1.0 Gyr.  We describe and analyse our models in Sect 2. Several Monte-Carlo simulations are presented and the results are shown in Sect 3. We will discuss the specific frequency of $N_{\\rm BS}$/$N_2$ of our simulations and compare our results to the observations in Sect 4. Summaries and conclusions are given in the final section. ",
        "conclusions": "\\label{sect:sum}  We investigated BSs generated from mass transfer scenario and simulated BS populations in young clusters in this work. BSs can be produced via various mechanisms (Piotto et al. 2004). Observed BSs in clusters can not be accounted for by a single mechanisms. In general, BSs could be formed by mass transfer in primordial binaries, which dominates BS formation in a spare environment (Mathys 1991), or through direct collisions, which are considered to be crucial in dense stellar environment (Fregeau et al. 2004).  The dynamical processes are not considered in this work because we would like to discuss those BSs only formed in mass transfer scenario. We choose young clusters as working samples in order to avoid the complicated effects of dynamical evolution. In reality, dynamic processes destroy primordial binaries while creating new ones through capture and exchange processes (Hurley et al. 2005). Mass segregation can also enhance BS formation in cluster cores. Magnetic activities, as discussed in Sect 2, can also affect the number of BSs through AML (Li et al. 2004; Demircan et al. 2006; Micheal \\& Kevin 2006; Stepien 2006), however it is only important in clusters older than $10^9$ yr (Chen \\& Han 2009).  We apply the same method as LDZ10 to calculate a grid of close binary evolution models undergoing both case A and B RLOF in a wider donor mass range. The formation and lifetime of BS are thoroughly analyzed using the models. Our results show that case B produces bluer, brighter and longer lived BSs than case A, therefore case B is more effective in producing BSs than case A, especially in young clusters.  Based on the grid of all our models at a given age, Monte-Carlo simulations were carried out for a number of clusters of AL07. The results show that BSs formed via different scenarios populate different areas on CMD. Mass transfer efficiency is considered in this work. A constant mass transfer efficiency of 0.5 is adopted for the comparison between conservative and non-conservative cases. The results show a high value of $\\beta$ (conservative case) can reproduce luminous BSs through case B scenario while lower value case ($\\beta$ = 0.5) can increase the number of BSs with low luminosity. Case A is less important in young clusters as predicted in Lu \\& Deng (2008). Case B and merger become more important for clusters with younger ages. The specific frequency of $f$ = $N_{\\rm BS}$/$N_2$ is studied. The result is less conclusive in the current scheme.  Two separated sequences can be noticed in CMD in the clusters between $\\log(age)>8.7$ up to nearly 1.0 Gyr (Fig~\\ref{f:bsn2}). Case B and merger models dominate the blue sequence, in which case B mainly populate the luminous part of the sequence. The case A models are only visible in a region close to the ZAMS in relatively older clusters. This may provide a new clue to the two sequences of BSs observed in star globular clusters (Ferraro et al. 2009).  The grid of detailed binary evolution models in this work covers a large parameter space, which can account for all possible binary mass transfer processes. Therefore, this grid of models can serve future studies on star clusters. By implementing dynamics, more realistic simulation of star clusters, including under populated open clusters, is possible. This will be part of our future work."
    },
    "1107/1107.4974_arXiv.txt": {
        "abstract": "{We present the analysis of two long, quasi-uninterrupted \\RXTE observations of Cygnus~X-1 that span several days within a 10\\,d interval. The spectral characteristics during this observation cover the region where previous observations have shown the source to be most dynamic. Despite that the source behavior on time scales of hours and days is remarkably similar to that on year time scales. This includes a variety of spectral/temporal correlations that previously had only been observed over \\Cyg's long-term evolution. Furthermore, we observe a full transition from a hard to a soft spectral state that occurs within less than 2.5\\,hours -- shorter than previously reported for any other similar \\Cyg transition.  We describe the spectra with a phenomenological model dominated by a broken power law, and we fit the X-ray variability power spectra with a combination of a cutoff power law and Lorentzian components. The spectral and timing properties are correlated: the power spectrum Lorentzian components have an energy-dependent amplitude, and their peak frequencies increase with photon spectral index. Averaged over 3.2--10\\,Hz, the time lag between the variability in the 4.5--5.7\\,keV and 9.5--15\\,keV bands increases with decreasing hardness when the variability is dominated by the Lorentzian components during the hard state. The lag is small when there is a large power law noise contribution, shortly after the transition to the soft state.  Interestingly, the soft state not only shows the shortest lags, but also the longest lags when the spectrum is at its softest and faintest.  We discuss our results in terms of emission models for black hole binaries.} ",
        "introduction": "\\label{sect:intro} The observed behavior of most black hole (BH) X-ray binaries (XRBs) can be classified into two major states, the so-called hard, power-law dominated, and the soft, thermally dominated state \\citep[e.g.,][and references therein; these authors also introduce the most extreme steep power law state seen in only a few BHs]{RemillardMcClintock2006}. In the hard state the X-ray continuum can be described by a hard power law, which is interpreted as resulting from Comptonization in a hot ($kT_\\mathrm{e}\\approx50$--100\\,keV) electron plasma surrounding the accretion disk, or in a jet \\citep{Dove1998,HaardtMaraschi1993,Zdziarski2002,markoff:05a}. A strong contribution of jet synchrotron emission is found in some sources \\citep{Markoff2001,Russell2010,Laurent2011}. In the case of \\Cyg the X-ray continuum can be modeled with a broken power law hardening above a break energy $E_\\text{break}\\approx10$\\,keV and a high-energy cutoff with a folding energy of $\\sim$100\\,keV \\citep[e.g.,][]{Wilms2006}. In the hard state, BH XRBs are generally observable in the radio with a flat/inverted radio through IR spectrum \\citep{FenderBelloniGallo2004,Fender2009,Dunn2010}. X-ray variability above 0.1\\,Hz during this state is high ($\\mathrm{rms}>20\\%$). In the soft state the X-ray spectrum is dominated by emission from an accretion disk and the core radio/IR emission is quenched. \\Cyg, however, retains a strong power law component even in its softest X-ray spectra.  In most Galactic BHs, transitions between states are common. The state behavior generally can be described as part of a characteristic evolution of spectral and timing states that depends on the previous source history. Over the outburst of an X-ray transient, the source is seen to move along a \\textsf{q}-shaped track in a hardness intensity diagram \\citep[HID; e.g.,][]{FenderBelloniGallo2004,HomanBelloni2005}. The outburst starts in the hard state; during the initial phase, the source brightens without significantly changing its spectral shape, and the radio/IR flux increases. At some point, the X-ray spectrum starts to soften. The radio/IR flux becomes more erratic.  In some sources radio ejections were seen during this intermediate state \\citep{Corbel2004,FenderBelloniGallo2004,Fender2006,Wilms2007}. After the state transition, the source spectrum is soft and no core radio/IR emission is detected \\citep{Fender2009}. During this soft state phase the bolometric source flux decreases, until at some lower luminosity the X-ray spectrum hardens, and the source becomes detectable in the radio/IR again. During the final part of the outburst the source is back in the hard state and the source luminosity decreases until the XRB ceases to be observable \\citep{FenderBelloniGallo2004,Belloni2006}.  The observed behavior of BH XRBs is determined by the interplay between emission from the accretion disk and from the regions of the accretion inflow/outflow in which the hard X-rays and the radio emission are generated. Transitions between the two states are therefore especially suited to place constraints on the physics of these components.  Despite the fact that the outbursts can last weeks or months, the transitions between hard and soft states are for most sources rare and often fast. \\citet{Belloni2006}, for example, report a hard to soft transition of \\mbox{GX\\,339$-$4} that occurred in less than 10\\,h. For this reason, the behavior during state transitions is not well studied. In addition, comparisons between different sources are complicated because that the influence of other variables, such as the variation of the mass accretion rate, is difficult to ascertain. Here we will present a detailed analysis of a state transition as observed in a single object.  The canonical BH candidate \\object{Cygnus X-1} is a persistent XRB that spends most of the time in the hard state; however, it shows frequent excursions into soft states \\citep[][and references therein]{Wilms2006}, including aborted transitions \\citep[``failed state transitions'';][]{Pottschmidt2000,Pottschmidt2003}.  Because of its brightness and persistent emission, \\Cyg is one of the best studied BHs.  In addition to studies of snapshot observations \\citep[e.g.,][]{gierlinski:99a,disalvo:01a,frontera:01a,makishima:08a,hanke:09a}, most recent work has concentrated on the long-term behavior of the source from a monitoring campaign with the \\textsl{Rossi X-Ray Timing Explorer} (\\RXTE) at a two week sampling that started in 1998 and is still ongoing \\citep[][and references therein]{Pottschmidt2003,gleissner:04a,gleissner:03a,Axelsson2005,Wilms2006}. The motivation of the present work was to fill the gap between the biweekly time scale and individual \\RXTE pointings with an intensive quasi-continuous observation spanning 10\\,days.  Our aim was to determine differences and similarities between the behavior of \\Cyg on time scales of minutes, hours, days, and the long-term behavior found during the monitoring campaign. As we will show in the following, this observation covered an especially interesting time when a state transition occurred.  The remainder of this paper is structured as follows: In Sect.~\\ref{sec:data} we describe the observations and data analysis methods, the results of which are presented in Sects.~\\ref{sec:analysis} and \\ref{sec:timing_analysis}. The results are discussed in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  \\subsection{Discussion}  \\subsubsection{State transition}  The changes of the emission properties of \\Cyg are continuous. Because the source does not follow the complete \\textsf{q}-shaped track in a hardness intensity diagram, which is found in outbursts of X-ray transients \\citep{FenderBelloniGallo2004,HomanBelloni2005,Dunn2010}, it is difficult to specify well-defined borders between its different emission states.  We find the following changes of properties: In the classical hard state of \\Cyg strong radio emission is observed, which is thought to originate from a steady jet. The intensity of the radio emission increases as the spectrum softens up to about $\\Gamma_1 \\approx 2.1$. Above this value, when \\Cyg leaves its classical hard state \\citep{Wilms2006}, the radio flux decreases in correlation with the hard X-ray flux \\citep{gleissner:04a}.  In the transitional state data analyzed here, we see an additional change of the emission properties at $\\Gamma_1 \\approx 2.5$, where the HID separates and the behavior of the time lags changes.  Remarkably, this result is indicated by both spectral and timing analysis. The data points fall into two clearly separated regions in the HID, also corresponding to different power spectral shapes. Transient BH XRBs only drastically change their timing properties during transitions from ``hard intermediate states'' to ``soft intermediate states''.  The fractional rms drops because of the disappearance of flat-top noise components (described as Lorentzians in this paper) in favor of a $1/f$ (power law) noise (Miyamoto et al., 1993; Takizawa et al., 1997; Nespoli et al., 2003; Homan \\& Belloni, 2005; Casella et al., 2004; Casella et al., 2005; Belloni et al., 2005; Belloni, 2009).  In addition, the ``soft intermediate state'' power spectra may contain specific QPOs referred to as type-A and type-B (Casella et al., 2005).  \\citet{Nespoli2003} found a QPO at about 6\\,Hz in an observation of GX\\,339$-$4 during a transition from the hard to the soft state. The frequency of the QPO was correlated to the PCA count rate. The power spectra on short time scales (16\\,s) of the data analyzed here did not show a comparable QPO, but it indicated that the peak frequencies of the Lorentzian profiles are correlated with the PCA count rate even on time scales shorter than one \\textsl{RXTE} orbit. During the hard state the PCA count rate is anticorrelated with the spectral hardness. Similar to \\Cyg, GX\\,339$-$4 shows increasing time lags as the soft count rate increases in the hard state and the flat-top noise disappears when the GX\\,339$-$4 softens into a ``soft intermediate state'' \\citep{Belloni2005}.  \\figHidLag  \\subsubsection{Noise components}  Our model consisting of only two Lorentzian profiles in addition to a power law with exponential cutoff is clearly justified by the structure of the power spectra shown in Fig.~\\ref{fig:g1_psd}. A third Lorentzian is only slightly indicated in a few of our hardest observations \\citep[consistent with][]{Pottschmidt2003}, but is not required in general.  Our identification of the Lorentzian profiles is equal to that of \\citet[see also the discussion in their Sect.~4.1]{Axelsson2005}, and our results are consistent with theirs: The $\\nu_1$--$\\nu_2$ relation is identical, and it is also visible here that rms$_2$ decreases uniformly with the peak frequency $\\nu_2$.  Power spectra in the two energy bands used here can be modeled simultaneously with the normalization as the only energy-dependent parameter. This means that a single physical process is responsible for producing the individual variability components, but that the contribution of each process to the variability is energy-dependent.  \\figCplRms  There are different physical models describing the origin and the behavior of the noise components observed in the power spectra. Some models connect the observed frequencies with fundamental general relativistic ones of perturbed orbits near the compact object \\citep[see][for a summary]{NowakLehr1998}. Among others, mechanisms discussed are perturbed epicyclic frequencies, the Lense-Thirring precession frequency, and ``diskoseismic'' modes that are related to maximum epicyclic frequency. Our observed peak frequencies, 0.5--30\\,Hz, correspond to epicyclic frequencies at about 22--345 gravitational radii for a $10M_\\odot$ compact object.  \\citet{NowakLehr1998} point out that the ``diskoseismic'' frequencies are very sensitive to luminosity changes, have a low rms variability, and thus cannot be applied to systems with $\\geq$10\\% rms as in our observation.  In addition, it is difficult to explain the spectral hardness of the variability components. Compared to the Lorentzian profile $L_1$, the one at high frequencies, $L_2$, is more dominant in the high-energy band, which is clearly above energies at which the disk emission contributes strongly, thus it is unlikely that both variabilities originate from the same region. \\citet{Wilkinson2009} observe a low-frequency Lorentzian profile originating from the accretion disk. \\citet{Uttley2011} point out that the disk can produce high-frequency variations as well, which would not show up as strong soft variations, provided that they occur at small disk radii so that only a small fraction of disk flux is modulated as they propagate inward. In contrast to $L_1$, the absolute variability contributions of $L_2$, and also that of the $1/f$ noise component, are comparable in both energy bands, however.  \\citet{Psaltis2000} postulate a transition radius in the inner disk acting as a low band-pass filter, which yields a dynamical model for the production of modulations, with a strong resonance to orbital, periastron- and nodal-precession. The model accounts for the observed PSDs and explains correlations between the characteristic frequencies. We find $\\nu_1 \\propto \\sqrt{\\nu_2}$, which cannot be identified clearly with one of the correlations between the modes calculated by \\citet{Psaltis2000}. An extrapolation of this relation implies a characteristic frequency of $\\sim$30\\,Hz, at which $\\nu_1 = \\nu_2$.  \\figGoneQuality  In the soft state the variability (rms) at high frequencies is lower, indicating a relatively steady state within the analyzed time scales. A possible interpretation of the power law arising in the PSDs as the spectrum softens is flicker noise in the accretion disk \\citep{Lyubarskii1997}. The data analyzed here indicate, however, that the variability contribution of the $1/f$ component is strong in both energy bands and slightly stronger in the high-energy band compared to the low-energy band in the cases of strong $1/f$ noise contribution, i.e., when the X-ray spectrum is soft (Fig.~\\ref{fig:cpl_rms}). A detailed interpretation of the soft state power spectra is presented by \\citet{Misra2008}, who consider linear and non-linear responses of a damped harmonic oscillator, which is assumed to govern the disk variability at each radius.  \\figRms  The dependence of the Lorentzians' rms on their peak frequency (Fig.~\\ref{fig:rms}) is similar to that presented by \\citet{Pottschmidt2003} in Figure 6 of their work. There is, however, one slight difference: Their Figure 6 indicates a unique relation between the peak frequency and the rms of all Lorentzians below $\\sim$20\\,Hz, i.e., when one Lorentzian shifted to a frequency at which another Lorentzian used to be, its rms was the same as the other Lorentzian had at this frequency. This relation could be interpreted as a self-similarity, i.e., one Lorentzian profile replaces another one when it shifts to its frequency and originates from the same process or the same mode. Figure~\\ref{fig:rms} does not show this behavior.  Owing to the different energy dependence of the Lorentzians it is obvious that this relation depends on the considered energy channels. It is nevertheless possible that one Lorentzian replaces another one at its frequency, because when L$_1$ shifts to the position of L$_2$ other properties of the source also change. For example $\\nu_2$ is about 5\\,Hz for $\\Gamma_1\\approx2.0$, whereas $\\nu_1$ reaches this frequency at $\\Gamma_1\\approx2.7$. This change in the X-ray spectrum implies changes in the emitting region, possibly affecting the properties of the variability components.  \\subsubsection{Time lags}  Previous time lag studies of \\Cyg showed typically short lags around 2--3\\,ms in the hard state, while enhanced lags up to 9\\,ms were observed during ``failed state transitions''. \\citet{Pottschmidt2003} used the energy bands 2--4\\,keV and 8--13\\,keV to obtain these values. When the source entered deeper into the soft state, the time lag returned to short values \\citep{Pottschmidt2000}.  Our new data set with higher temporal resolution (\\RXTE orbit-wise extractions) allows us to extend these results (Fig.~\\ref{fig:multi_timing}, \\ref{fig:hid_lag}). In the hard state the time lag increases with softening X-ray spectrum consistent with previous observations, whereas after the transition to the soft state the time lag behavior changes. Depending on the brightness, the time lag changes from values of up to 11\\,ms at low count rates to values of 2--3\\,ms at high count rates without appreciable softening of the source. Changes from long to short time lags occur from one \\RXTE orbit to the next. Figure~\\ref{fig:multi_timing} shows that the lag is short when the variability contribution of the power law is large, whereas the time lags of the Lorentzian variability are long depending on their peak frequency and thus on the spectral hardness.  Time lags can originate from reprocessing of radiation and propagation in the emitting region. \\citet{Nowak1999compton} exclude disk-intrinsic time delays as the sole reason for the lag, and discuss constraints on properties of the corona if the time lag originates in it. \\citet{Kylafis2008} present a jet model that can reproduce the behavior of the time lags in the \\Cyg hard state. \\citet{KoerdingFalcke2004} show that time lags can be caused by a pivoting power law spectrum arising, e.g., from a jet/synchrotron model. In their work the sign of the time lag depends on the hardness/flux correlation \\citep[][Fig.~9]{KoerdingFalcke2004}. Interestingly, the correlation visible in the lambda-shaped HIDs seems to change its sign coinciding with the change in time lag behavior, exactly at the transition: in the hard state the count rate increases with decreasing hardness, while it is the other way around in the SIMS. The transition is also where the behavior of the time lag changes.  \\citet{Uttley2011} show that the time lags seen at soft X-rays (0.5--9\\,keV) are caused by the accretion disk. They point out that Comptonization models for lags have difficulties, whereas propagation delays \\citep[see, e.g.,][]{Kotov2001,Arevalo2006} can explain their data, and that a similar mechanism can cause time lags at higher energies.  The clear correlations between spectral and timing properties put strong constraints on models explaining the emission of accreting black holes.  \\subsection{Summary} The main results of this paper can be summarized as follows: \\begin{itemize} \\item \\Cyg was in an intermediate/soft state during 2005 February 01--11. The photon index $\\Gamma_1$ reached values between 2.0 and 2.7 and was correlated with $\\Gamma_2$ in the same way as found by \\citet{Wilms2006}. \\item The source was very variable during that period of time; the total PCA count rate varied between 1000 and 5000 counts$\\rm\\:s^{-1}$ per PCU. \\item A full transition from the hard state to the soft state occurred in less than 2.5\\,hours. The two states can be clearly distinguished in a hardness intensity diagram. \\item Spectral and timing parameters are strongly correlated. The peak frequencies of the Lorentzian profiles shift to higher values with increasing $\\Gamma_1$. Those are identical in the low- and the high-energy band, whereas their rms depends on the energy. \\item In the HIMS the time lag increases with decreasing hardness. After the transition to the SIMS the time lag behaves differently. At high count rates (${\\gtrsim}2500\\,\\mathrm{counts}\\:\\mathrm{s}^{-1}\\,\\mathrm{PCU}^{-1}$ in this case) the time lag drops to about 2\\,ms and the contribution of the power law noise is large. At lower count rates the time lag is long, with values up to 11\\,ms. \\end{itemize}"
    },
    "1107/1107.3553_arXiv.txt": {
        "abstract": "The atomic CII fine-structure line is one of the brightest lines in a typical star-forming galaxy spectrum with a luminosity $\\sim 0.1$\\% to 1\\% of the bolometric luminosity. It is potentially a reliable tracer of the dense gas  distribution at high redshifts and could provide an additional probe to the era of reionization. By taking into account of the spontaneous, stimulated and collisional emission of the CII line, we calculate the spin temperature and the mean intensity as a function of the redshift. When averaged over a cosmologically large volume,  we find that the CII emission from ionized carbon in individual galaxies is larger than the signal generated by carbon in the intergalactic medium (IGM). Assuming that the CII luminosity is proportional to the carbon mass in dark matter halos, we also compute the power spectrum of the CII line intensity at various redshifts. In order to avoid the contamination from CO rotational lines at low redshift when targeting a CII survey at high redshifts, we propose the cross-correlation of CII and 21-cm line emission from high redshifts. To explore the detectability of the CII signal from reionization, we also evaluate the expected errors on the CII power spectrum and CII-21 cm cross power spectrum based on the design of the future milimeter surveys. We note that the CII-21 cm cross power spectrum contains interesting features that captures physics during reionization, including the ionized bubble sizes and the mean ionization fraction, which are challenging to measure from 21-cm data alone. We propose an instrumental concept for the reionization CII experiment  targeting the frequency range of $\\sim$ 200 to 300 GHz with 1, 3 and 10 meter apertures and a bolometric spectrometer array with 64 independent spectral pixels with about 20,000 bolometers. ",
        "introduction": "Carbon is one of the most abundant elements in the Universe and it is singly ionized (CII) at $11.26$ eV, an ionization energy that is less than that of the hydrogen. With a splitting of  the fine-structure level at $91\\ \\rm K$ CII is easily excited resulting in a line emission at $157.7\\ \\rm \\mu m$ through the $^2P_{3/2}\\to ^2P_{1/2}$ transition. It is now well established that this line provides a major cooling mechanism for the neutral interstellar medium (ISM) \\citep{Dalgarno72, Tielens85,Wolfire95,Lehner04}. It is present in multiple phases of the ISM in the Galaxy \\cite{Wright91} including the most diffuse regions \\cite{Bock93} and the line emission has been detected from the photo dissociation regions (PDRs) of star-forming galaxies \\citep{Boselli02,DeLooze11,Nagamine06,Stacey10} and, in some cases, even in $z > 6$ Sloan quasars \\citep{Walter09}.  The CII line is generally the  brightest emission line in star-forming galaxy spectra and contributes to about 0.1\\% to 1\\% of the total far-infrared (FIR) luminosity \\citep{Crawford85,Stacey91}. Since carbon is naturally produced in stars, CII emission is then expected to be a good tracer of the gas distribution in galaxies. Even if the angular resolution to resolve the CII emission from individual galaxies is not available, the brightness variations of the CII line intensity can be used to map the underlying distribution of galaxies and dark matter \\citep{Basu04,Visbal10,Gong11}.  Here we propose CII intensity mapping as an  alternative avenue to probe the era of reionization, including the transition from primordial galaxies with PopIII stars alone to star-formation in  the second-generation of galaxies with an ISM polluted by  metals. CII intensity mapping complements attempts to study reionization with low-frequency radio experiments that probes the 21-cm spin-flip line from neutral hydrogen. While those experiments are sensitive to the neutral gas distribution dominated by the IGM, CII will probe the onset of star formation and metal production in $z \\sim 6$ to 8 galaxies.  Recently it has also been proposed to use rotational lines of CO molecule to probe reionization (e.g., Gong et al. 2011; Carilli 2011; Lidz et al. 2011). CO studies have the advantage that redshift identification is facilitated by multiple $J$ transition lines, while with CII some confusion could result in the line identification with other atomic fine-structure and molecular lines at sub-mm and mm wavelengths. In comparison, and based on limited observational data, CII emission is expected to brighter in low mass galaxies, compared to the case of CO luminosity. It is unclear if the galaxies present during reionization are analogous to low-redshift dwarf galaxies or the local ultra-luminous infrared galaxy population. At high redshifts, we expect most of the carbon to be in CII  rather than CO, since the high-redshift galaxies may not host significant dust columns required to shield CO from dissociating UV photos. Given that a CO experiment to study reionization will involve an experiment at mid radio frequencies of 15 to 30 GHz, while CII will involve high frequencies, the two methods will be affected by different systematics and foregrounds. Thus, a combined approach involving multiple probes of reionization, 21-cm, CO and CII, would provide the best avenue to study the $z > 6$ universe.  In this paper we present an analytical calculation to predict the CII intensity intensity as a function of the redshift by considering spontaneous, stimulated and collisional emission processes \\citep{Suginohara99,Basu04}. This intensity changes the brightness spectrum of the cosmic microwave background (CMB) at the frequency corresponding to the CII line (Basu et al. 2004).  In this paper, we focus on the CII flux from individual galaxies where the matter density is high and the collisional emission is the dominant process. As a check on our analytical calculations, we also consider results derived from numerical simulations to establish the CII intensity. The two approaches are generally consistent. We then consider the measurement of the CII intensity fluctuations resulting from the clustering of the galaxy distribution and sources that are present during and towards the end of reionization at $z \\sim 6-8$.  Experimentally, there are several challenges that one must overcome before the CII intensity fluctuations from high redshifts can be reliably established. First, higher J transitions of CO from dense molecular gas at lower redshifts contaminate the CII line intensity measurements. In particular, one must account for all of CO(2-1) to CO(13-12) emission lines from individual narrow redshift ranges in the foreground between 0 and 6 when probing CII fluctuations at $z > 6$. To the extent to which a variety of existing predictions on the CO intensity can be trusted \\citep{Gong11,Lidzetal11,Carilli11}, we generally find that the contamination is mostly below the level of the CII signal. Extending previous studies \\citep{Basu04,Visbal10,Gong11,Lidzetal11}, we propose the cross-correlation of CII line intensity mapping and 21-cm fluctuations as a way to improve studies related to the epoch of reionization. To evaluate the detectability of the CII signal, we calculate the errors on the CII power spectra and CII-21cm cross correlation, respectively, based on the design of potential (sub-)milimeter  surveys for the CII emission. For 21-cm, we consider the first generation experiment LOw Frequency ARray\\footnote{http://www.lofar.org/} (LOFAR) survey as well as the improvement expected from a second generation experiment like the low-frequency extension to the Square Kilometer Array (SKA).  The paper is organized as follows: in the next section, we derive the formulas to calculate the spin temperature of the CII line in the ISM of galaxies and the IGM. In Section~3, we calculate the mean CII intensity analytically and compare it with results derived from a simulation. We show the CII power spectrum in Section~4. In Section~5, we discuss the low-redshift contamination from CO emission lines for CII intensity mapping at $z > 6$ and in Section~6 we propose a  cross-correlation between CII and 21-cm line intensity measurements over the overlapping redshift ranges as a way to both distinguish the CII signal from the CO emission and to improve overall constraints on reionization physics. In Section~7 we outline the experimental parameters of a CII spectral mapping experiment designed to probe $z \\sim 6$ to 8 CII intensity fluctuations and discuss the detectability of the CII power spectrum and the CII and 21-cm cross-correlation when combined with LOFAR and SKA. We conclude with a summary of our results in Section~8. Throughout this paper, we assume the flat $\\rm \\Lambda CDM$ model with $\\Omega_{m} = 0.27$, $\\Omega_b=0.046$, $\\sigma_8=0.81$, $n_s=0.96$ and $h=0.71$ \\citep{WMAP7}. ",
        "conclusions": "In this paper, we have estimated the mean intensity and the spatial intensity power spectrum of the  CII emission from galaxies at $z > 6$. We first calculated the CII intensity analytically for both the ISM of galaxies and the diffuse IGM  and showed that the CII emission from dense gas in the ISM of galaxies is much stronger than that from the IGM. Then, to check the analytical calculation, we used the hot gas or the hot+warm gas from a simulation to find the CII mass in a halo, and further calculated the CII number density and intensity. We found that the two methods are in good agreement especially at high redshifts we are interested in. Next, we computed the CII clustering power spectrum assuming the CII luminosity is proportional to the CII mass, $L_{\\rm CII}\\sim M_{\\rm CII}$. We compared our CII power spectrum with that derived from the $L_{\\rm CII}$-$L_{\\rm CO(1-0)}$ relation, and found that they are consistent in the $1\\sigma$ level. We also explored the contamination of the CII emission by the CO lines at lower redshift, and found the contamination can lead to $2\\%$ and $30\\%$ enhancement of the CII power spectrum at $z=7$ and $z=8$.  To reduce the foreground contamination and to improve our scientific understanding of reionization we propose here a cross-correlation study between the CII and 21-cm emission in the overlapping redshift ranges and the same part of the sky. The cross-correlation exists since they both trace the same matter distribution. At large scales the correlation is due to ionized bubbles surrounding CII bright galaxies while at smaller scales both CII galaxies and ionized bubbles trace the underlying density field. We have outlined 3 potential CII experiments using 1, 3 and 10 m aperture telescope outfitted with a bolometer array spectrometer with 64 independent spectral pixels. A 1 or 3 meter aperture CII experiment is matched to a first generation 21-cm experiment such as LOFAR while an improved CII experiment can be optimized to match a second generation 21-cm experiment like SKA. We find that the overall ability to extract details on reionization requires a careful coordination and an optimization of both CII and 21-cm experiments. We have not discussed the issues related to foregrounds, including AGN-dominated radio point sources that dominate the low-frequency observations and dusty galaxies that dominate the high-frequency CII observations and Galactic foregrounds such as dust and synchrotron. In future papers we will return to these topics and also expand our discussion related to the instrumental concept as we improve the existing outline to an actual design."
    },
    "1107/1107.2249_arXiv.txt": {
        "abstract": "{Abstract: The Canada-France-Hawaii Legacy Survey (CFHTLS) comprising about 25,000 MegaCam images was data mined to search for serendipitous encounters of known Near Earth Asteroids (NEAs) and Potentially Hazardous Asteroids (PHAs). A total of 143 asteroids (109 NEAs and 34 PHAs) were found on 508 candidate images which were field corrected and measured carefully, and their astrometry was reported to Minor Planet Centre. Both recoveries and precoveries (apparitions before discovery) were reported, including data for 27 precovered asteroids (20 NEAs and 7 PHAs) and 116 recovered asteroids (89 NEAs and 27 PHAs). Our data prolonged arcs for 41 orbits at first or last opposition, refined 35 orbits by fitting data taken at one new opposition, recovered 6 NEAs at their second opposition and allowed us to ameliorate most orbits and their Minimal Orbital Intersection Distance (MOID), an important parameter to monitor for potential Earth impact hazard in the future. } ",
        "introduction": "\\label{intro}  Despite the continuous grow of the existing imaging archives and surveys taken with various telescopes around the globe, extremely little work has been devoted to data mining in order to ameliorate the orbits of known asteroids and Near Earth Asteroids (NEAs) and Potentially Hazardous Asteroids (PHAs).  Searching for known minor bodies in old imaging archives is not a new idea, such work being carried in the last two decades by a few authors in order to recover some asteroids and comets and improve their orbits (\\cite{bow92}; \\cite{hav92}; \\cite{mcn95}; \\cite{bf00}; etc). During the last decade, some dedicated data mining work has been carried out to search for known NEAs in a few entire photographic plate archives, namely the projects AANEAS (\\cite{ste98}, who introduced the term ``precovery''), ANEOPP (\\cite{boa01}) and DANEOPS (\\cite{hah02}). Recently, we presented the public server PRECOVERY devoted to search {\\it all} known asteroids (including NEAs and PHAs besides all other catalogued asteroids) in {\\it any} archive uploaded by the user, given by a simple observing log recorded in a standard format (\\cite{vad09}).  More than 7,400 NEAs are known today (Nov 2010) and some 1,170 of these are catalogued as PHAs according to the JPL NEO database (\\cite{nas10}). Many of these bodies have been insufficiently observed, i.e. only during about one month of visibility at their first opposition. Some of these are classified as Virtual Impactors (VIs), while about 70 are considered lost due to their present very large uncertainty in their orbits and ephemeris and their very faint brightness, according to the NEODyS database (\\cite{mil10a}). Based on the currently available observations, JPL Sentry System (\\cite{nas10}) monitors more than 300 NEAs possibly to cause future Earth impact events during the next 100 years, although virtually all have almost zero probability to cause such impacts.  Thanks to five dedicated US-lead surveys searching for NEAs during the last two decades, we have discovered the tip of the iceberg of the entire NEA population consisting mostly in $\\sim1$km and larger objects detectable with 1m class telescopes. Nevertheless, sub-km sized asteroids as small as 150m could still cause regional or global scale disaster in the eventuality of a catastrophic event (\\cite{mor06}), thus a common effort should be pursued not only to discover but also to recover and follow-up known NEAs.  During the past 4 years, part of the EURONEAR project we have observed about 200 selected NEAs using 10 non-dedicated 1-2m telescopes during about 50 nights obtained mostly through regular time allocation competition which has been difficult to obtain in the absence of a dedicated facility (\\cite{bir10}). Besides new observations, data mining of existing imaging archives represents another goal of the EURONEAR program, and a first paper introduced the method and software to perform the search of any archive for known NEAs, PHAs and other asteroids (\\cite{vad09}).  In the present paper we will use the same method to data mine the entire Canada-France-Hawaii Legacy Survey (more than 25,000 wide field MegaCam images) for known Near Earth Asteroids. Section~\\ref{CFHTLS} briefly introduce the survey and present the data mining method. Section~\\ref{results} will present the results grouped in five special classes, and Section~\\ref{future} will conclude the paper, introducing two related projects in development. ",
        "conclusions": ""
    },
    "1107/1107.3659_arXiv.txt": {
        "abstract": "We constrain the evolution of Newton's constant using the growth rate of large-scale structure measured by the WiggleZ Dark Energy Survey in the redshift range $0.1 < z < 0.9$.  We use this data in two ways. Firstly we constrain the matter density of the Universe, $\\omms$ (assuming General Relativity), and use this to construct a diagnostic to detect the presence of an evolving Newton's constant. Secondly we directly measure the evolution of Newton's constant, $\\Geff$, that appears in Modified Gravity theories, without assuming General Relativity to be true. The novelty of these approaches are that, contrary to other methods, they do not require knowledge of the expansion history of the Universe, $H(z)$, making them model independent tests. Our constraints for the second derivative of Newton's constant at the present day, assuming it is slowly evolving as suggested by Big Bang Nucleosynthesis constraints, using the WiggleZ data is $\\ddotGeff(t_0)=-1.19\\pm 0.95\\cdot 10^{-20}h^2 \\textrm{yr}^{-2}$, where $h$ is defined via $H_0=100~h~km~ s^{-1}Mpc^{-1}$, while using both the WiggleZ and the Sloan Digital Sky Survey Luminous Red Galaxy (SDSS LRG) data is $\\ddotGeff(t_0)=-3.6\\pm 6.8\\cdot 10^{-21}h^2 \\rm{yr}^{-2}$, both being consistent with General Relativity. Finally, our constraint for the rms mass fluctuation $\\sigma_8$ using the WiggleZ data is $\\sigma_8=0.75 \\pm 0.08$, while using both the WiggleZ and the SDSS LRG data $\\sigma_8=0.77 \\pm 0.07$, both in good agreement with the latest measurements from the Cosmic Microwave Background radiation. ",
        "introduction": "Recent observations suggest that the Universe is undergoing a phase of accelerated expansion, usually attributed to an unknown ideal fluid dubbed dark energy. Most of the evidence for the existence of dark energy comes from geometric tests that measure the expansion rate of the universe $H(a)\\equiv \\dot{a}/a$ at different epochs, where $a(t)$ is the scale factor of a Friedmann-Robertson-Walker metric. Examples of such tests include measurements of the luminosity distance, $d_L(a)$, using standard candles like Type Ia supernovae (SN Ia) \\cite{Kilbinger:2008gk, Kessler:2009ys} and measurements of the angular diameter distance, $d_A(a)$, using standard rulers such as the scale of the sound horizon at last scattering \\cite{Komatsu:2010fb} and baryon acoustic oscillations \\cite{Percival:2009xn}. Even though these tests are presently the most accurate probes of dark energy, the mere determination of the expansion rate, $H(a)$, is not able to provide significant insight into the properties of dark energy to distinguish it from models that attribute the accelerating expansion to modifications of general relativity \\cite{Tsujikawa:2010zza,Heavens:2007ka,Nesseris:2006er,Nesseris:2006jc}.  A smoking gun signature of modified gravity theories, such as $f(R)$ gravity  \\cite{fRpapers} or scalar-tensor models \\cite{stensor}, is the fact that they predict that Newton's constant $G_N$ is evolving with time. Generically, modified gravity theories can be cast in such a way as to have an effective Newton's constant that changes with time $\\Geff(t)$ \\cite{DeFelice:2010gb, Tsujikawa:2007gd, DeFelice:2010aj}. However, this variation cannot be arbitrarily large and has already been constrained by Big Bang Nucleosynthesis to be within $10\\%$.  See, for example, Ref.~\\cite{Bambi:2005fi} which gives $\\Geff(t_{BBN})/\\Geff(t_0)=1.09\\pm^{0.22}_{0.19}$, where $t_{\\rm BBN}$ is the time of Big Bang Nucleosynthesis, while $t_0$ is the present day.  This allows us to make the simplifying assumption that $\\Geff$ is a slowly varying function. Furthermore, Solar System tests and SN Ia observations (see \\cite{Nesseris:2006hp} and references therein) suggest that the first time derivative $\\dotGeff(t_0)$ is almost zero and that for the second time derivative we have $|g_2|\\lesssim O(1)$ , where $ g_2 \\equiv \\ddotGeff(t_0)/(\\Geff(t_0)H_0^2)$.  The additional observational input that is required to discriminate between dark energy and modified gravity theories is the growth function of the linear matter density contrast $\\delta\\equiv\\frac{\\delta\\rho}{\\rho}$, where $\\rho$ represents the background matter density and $\\delta\\rho$ its first order perturbation. The reason for this is the fact that the growth of cosmic structure is a result of the motion of matter and therefore is sensitive to both the expansion of the Universe $H(a)$ and $\\Geff$.  More specifically, it can be shown that in modified gravity theories on sub-horizon scales, ie $k^2\\gg a^2H^2$ where $k$ is the wave-number of the modes of the perturbations in Fourier space, the growth factor satisfies the following differential equation \\cite{DeFelice:2010gb, Tsujikawa:2007gd, DeFelice:2010aj, Nesseris:2009jf}: \\be \\delta''(a)+\\left(\\frac{3}{a}+\\frac{H'(a)}{H(a)}\\right)\\delta'(a) -\\frac{3}{2}\\frac{\\omms \\Geff(a)}{a^5 H(a)^2/H_0^2}~\\delta(a)=0 \\label{ode}\\ee where primes denote differentiation with respect to the scale factor,  $H(a)\\equiv\\frac{\\dot{a}}{a}$ is the Hubble parameter and we assume the initial conditions $\\delta(0)=0$ and $\\delta'(0)=1$ for the growing mode. When $\\Geff(a)=1$ we get GR as a subcase, while in general in modified gravity theories $\\Geff(a)$ is time and scale dependent. For example, in $f(R)$ theories we have that \\cite{Tsujikawa:2007gd} \\ba \\Geff&=&\\frac{1}{8\\pi F}\\frac{1+4\\frac{k^2}{a^2}m}{1+3\\frac{k^2}{a^2}m} \\label{gefffr1}\\\\ m&\\equiv& \\frac{F_{,R}}{F}\\\\F&\\equiv&f_{,R}=\\frac{\\partial f}{\\partial R}\\label{gefffr2}\\ea which reduces to GR only when $f(R)=R-2\\Lambda$.  The exact solution of Eq.~(\\ref{ode}) for a flat GR model with a constant dark energy equation of state $w$  is given for the growing mode by \\cite{Silveira:1994yq,Percival:2005vm}:\\ba \\delta(a)&=& a \\cdot {}_2F_1 \\left(- \\frac{1}{3 w},\\frac{1}{2} - \\frac{1}{2 w};1 - \\frac{5}{6 w};a^{-3 w}(1 - \\omms^{-1})\\right) \\label{Da1} \\nn \\\\ \\textrm{for}&&H(a)^2/H_0^2= \\omms a^{-3}+(1-\\omms)a^{-3(1+w)},\\ea where ${}_2F_1(a,b;c;z)$ is a hypergeometric function defined by the series \\be {}_2F_1(a,b;c;z)\\equiv \\frac{\\Gamma(c)}{\\Gamma(a)\\Gamma(b)}\\sum^{\\infty}_{n=0}\\frac{\\Gamma(a+n)\\Gamma(b+n)}{\\Gamma(c+n)n!}z^n \\ee on the disk $|z|<1$ and by analytic continuation elsewhere, see Ref.~\\cite{handbook} for more details.  In more general cases it is not possible to find an analytical solution for Eq.~(\\ref{ode}), but in Ref.~\\cite{Wang:1998gt} it was shown that the growth rate $f(a)\\equiv \\frac{d ln \\delta}{dlna}$ can be approximated as \\ba f(a)=\\omms(a)^\\gamma \\label{fg}\\\\\\omms(a)\\equiv\\frac{\\omms~a^{-3}}{H(a)^2/H_0^2} \\\\ \\gamma \\simeq \\frac{3 (1-w)}{5-6 w} \\ea where this approximation for $\\gamma$ is  valid at first order for a dark energy model with a constant equation of state $w$ and for \\lcdm we have $\\gamma=\\frac{6}{11}\\simeq 0.55$. Conveniently, the approximation of Eq.~(\\ref{fg}) can be used to accurately fit a wide variety of different scenarios, including modified gravity models, if $\\gamma$ is allowed to be a free parameter.  In order to get a prediction for $\\Geff$ from the growth rate data one would need to solve Eq.~(\\ref{ode}) and this can only be done for a specific theory, as knowledge of $H(a)$ is required. Therefore, in order to reconstruct $\\Geff$ the usual procedure would be to assume a model for $H(a)$, fit the approximation of Eq.~(\\ref{fg}) to the data and finally use Eq.~(\\ref{ode}) to get a prediction for $\\Geff$. This obviously has several drawbacks, as the growth data alone are not sufficient to get viable fits to $H(a)$ and are usually complemented by SN Ia and CMB data. The SN Ia data actually determine the luminosity distance $d_L(a)$, therefore converting these measurements to H(a) depends on differentiation of noisy data and an assumption about the spatial curvature $\\Omega_K$ and this obviously results in large error bars for the best fit parameters. Finally, any result would ultimately depend on the specific choice of $H(a)$, or in other words it would suffer from model-choice bias, whereas what we propose is a model independent technique that can be used even without the CMB and SN Ia data. In order to solve these problems we constructed a diagnostic that is based solely on data, with no assumptions about the expansion history of the Universe, that is able to detect the presence of an evolving Newton's constant. We describe both our new method and the diagnostic in detail in Section \\ref{method}. ",
        "conclusions": "We presented a method that can be used to determine the value of the second derivative of Newton's constant $\\ddotGeff(t_0)$, under the assumption of a slowly varying $\\Geff(a)$, by using the growth rate data of the WiggleZ survey. The novelty of our approach lies in the fact that contrary to other methods, this one does not require knowledge of the expansion history of the Universe $H(z)$, usually found separately by using the Type Ia Supernovae data but instead we can use the generic approximation of Eq.~(\\ref{fg}) which, as we have shown in the previous section, is able to fit a great variety of different scenarios (evolving equation of state and/or an evolving Newton's constant).  In order to demonstrate the ability of our method to measure the correct values of $g_2$ we performed simulations with synthetic noise-free and noisy data. We used four ``fiducial\" models, the analytic solution to Eq.~(\\ref{ode}) $M_1$, the generic approximation of Eq.~(\\ref{fg}) with the parameter $\\gamma$ fixed $M_2$ or free to vary $M_3$ and a polynomial expansion $M_4$. We tested the models with data created for different parameters $(w_0,w_a)$ of a dark energy equation of state $w(a)=w_0+w_a(1-a)$, three different values of the second derivative of Newton's constant $g_2 \\equiv \\ddotGeff(t_0)/(\\Geff(t_0)H_0^2)=(-1,0,1)$.  All models proved to be able to measure the correct value of $g_2$ within $1\\sigma$ of its true value, despite the fact that we tested them against data created from an evolving equation of state $w(z)$. The most accurate model, ie the one that was closer to the ``true\" value of $g_2$ used to create the synthetic data, was $M_3$ which gave an exact value when $g_2=-1$ and was roughly within $10\\%$ and $20\\%$ in the $g_2=0$ and $g_2=1$ respectively. The most precise model, ie the one with the smaller error bars, was found to be $M_2$ for $g_2=0$ and $g_2=-1$, but $M_1$ for $g_2=1$. So, we found that both the simulations with synthetic noise-free and noisy data lead to the conclusion that the model $M_3$ provides the most accurate measurements of the value of $g_2$ while having a precision comparably good to the other models.  When we applied our method to the real WiggleZ data, we found that the models $M_1$, $M_2$, $M_3$ and $M_{3, w=-1}$ give on average a value $g_2\\simeq-1.14 \\pm 0.91$ which corresponds to a value for the second derivative of $\\ddotGeff(t_0)=-1.19\\pm 0.95\\cdot 10^{-20}h^2 \\textrm{yr}^{-2}$. Also, their prediction for the matter density ($\\ommgr=0.40\\pm0.10$) is in agreement with the one measured from the WMAP probe ($\\omms=0.27\\pm0.03$) \\cite{Komatsu:2010fb}. Finally, we found that the polynomial $M_4$ gives a measurement of $\\sigma_8=0.75 \\pm 0.08$, which is in agreement with the corresponding WMAP 7 value $\\sigma_8=0.80 \\pm 0.03$ \\cite{Komatsu:2010fb}.  When we performed a joint analysis of the WiggleZ and the SDSS-II data of Ref.~\\cite{Samushia:2011cs}, the models $M_1$, $M_2$, $M_3$ and $M_{3, w=-1}$ gave on average a value $g_2\\simeq-0.34 \\pm 0.65$ which corresponds to a value for the second derivative of $\\ddotGeff(t_0)=-3.6\\pm 6.8\\cdot 10^{-21}h^2 \\textrm{yr}^{-2}$ and is again compatible with GR. Finally, the polynomial $M_4$ gives a measurement of $\\sigma_8=0.77 \\pm 0.07$, which is in agreement with the corresponding WMAP 7 value $\\sigma_8=0.80 \\pm 0.03$ \\cite{Komatsu:2010fb}.  We also presented a diagnostic \\be {\\O}\\equiv \\frac{\\omms-\\Omega_{\\rm m,GR}}{\\omms}=1-\\frac{\\Omega_{\\rm m,GR}}{\\omms}=1-\\frac{1}{3 I_0 \\omms} \\nn\\ee where $\\omms$ is the value of the matter density as measured from other independent observations. As is easily seen the value of $\\O$ is zero only in GR and non-zero in general in the presence of an evolving $\\Geff$. Furthermore, we found that the sign of $\\O$ is correlated to the sign of the second derivative $\\ddotGeff$, ie they are both either positive or negative. Therefore, if we assume that the value of $\\omms$ is independently and accurately determined by other observations, then any statistically significant deviation of the quantity ${\\O}$ from zero, clearly and uniquely identifies the presence of an evolving $\\Geff$ and consequently modified gravity theories. We should note that while the estimate of $g_2$ depends on the assumption of a slowly varying $\\Geff$ on the form of Eq.~(\\ref{geff}), the estimate of $\\O$ is model independent of any assumption on either  $\\Geff$ or the dark energy equation of state $w(a)$. Applying it to the real WiggleZ data we found $\\O=-0.486\\pm0.369$, which is consistent with zero.  Finally, we also determined the optimal redshift range where new data points should be attempted to be measured in future surveys. We noticed that due to degeneracies in $\\fs(a)$ there are sour-spots at high redshifts, ie there points where $\\fs(a)$ is the same regardless of the values of the parameters $w$ or $g_2$. So, in terms of adding new data points, we found that these should be focused in the range $0 < z < 0.8$.  In order to test whether having more points at low redshifts can help in the detection of an evolving $\\Geff$ and to estimate the necessary number of new points necessary to achieve this, we created synthetic data with noise following the same procedure as in Section IV. We placed 4 of the points fixed at the WiggleZ redshifts $z=(0.22, 0.41, 0.60, 0.78)$ and we also added 1, 2, 4, 8, 12 or 14 points randomly placed in three redshift ranges $0<z<0.8$, $1<z<1.4$ and $1<z<1.8$. Having more data points in any of the tested redshift ranges of $z>1$ allowed a maximum 1.7$\\sigma$ detection of $g_2=-1$, while more sticking to the lower redshift range of $0<z<0.8$ allowed a much more significant detection, at 3$\\sigma$, of $g_2=-1$.  In conclusion, we presented a method by which growth of structure could detect a variation in $\\Geff$ in a model-independent manner, but found that the current growth rate data is consistent with a non-evolving value of $\\Geff$. It is therefore imperative to have both more and better data points in order to possibly detect an evolving $\\Geff$ with any statistical significance. This makes it tantalizing to pursue this analysis further as both the quality and the number of the growth rate data will increase in the near future."
    },
    "1107/1107.1376_arXiv.txt": {
        "abstract": "We present  the final results from our ultra-deep spectroscopic campaign with FORS2 at the ESO/VLT for the confirmation of $z\\simeq7$ ``z--band dropout'' candidates selected from our VLT/Hawk-I imaging survey over three independent fields. In particular we report on  two newly discovered galaxies at redshift $\\sim$ 6.7 in the NTT deep field: both galaxies show a \\lya emission line with rest-frame EWs of the order 15-20\\AA\\ and luminosities of 2-4$\\times 10^{42} erg s^{-1}$. We also present the results of ultra-deep observations of a sample of i-dropout galaxies, from which we set a solid upper limit on the fraction of interlopers.  Out of the 20 z-dropouts observed we confirm 5 galaxies at $6.7 < z < 7.1$. This is systematically  below the expectations drawn on the basis of lower redshift observations: in particular there is a significant lack of objects with intermediate \\lya EWs (between 20 and 55 \\AA). We conclude that the trend for the fraction of  \\lya emission in LBGs  that is constantly increasing  from z$\\sim$3 to z$\\sim$6 is most probably reversed from $z\\sim 6$ to  z$\\sim$7. Explaining the observed rapid change in the LAE fraction among the drop-out population with reionization requires a fast evolution of the neutral fraction of hydrogen in the Universe. Assuming that the Universe is completely ionized at z=6 and adopting  the semi-analytical models of Dijkstra et al. (2011), we find that our data require a change of the neutral hydrogen fraction of the order $\\Delta\\chi_{HI} \\sim 0.6$ in a time $\\Delta z\\sim1$, provided that the escape fraction does not increase dramatically over the same redshift interval. ",
        "introduction": "The epoch of reionization marks a major phase transition of the Universe, during which the intergalactic space became transparent to UV photons. Determining when this occurred and the physical processes involved represents the latest frontier in observational cosmology. Over the last few years, searches have intensified to identify the population of high-redshift galaxies that might be responsible for this process. Recent data  from the Hubble Space Telescope WFC3 infrared camera provided a dramatic advance accessing  the faint side of the z$\\sim 7 $ UV Luminosity Function (LF) \\citep{McLure2010,Bouwens2010a,Grazian2011}, while wide field ground-based surveys are starting to detect the brightest galaxies at z$\\sim 7$ \\citep{Ouchi2009,Castellano2010a,Castellano2010b}. These surveys provided evidence that the UV LF evolves significantly with redshift \\citep{Wilkins2010}, that the UV continuum of galaxies appears progressively bluer with increasing z (\\cite{Fin2010,Bouwens2010b} but see also \\cite{Dunlop2011}) and that their stellar masses are on average smaller than those of their lower-z counterparts \\citep{Labbe2010}, while the morphologies remain essentially unchanged from z $\\sim$6 to z$\\sim$7 \\citep{Oesch2010}.  Unfortunately, these results are entirely based on color selected samples.  The lack of spectroscopic redshifts prevents us from deriving firm conclusions based on these early samples.  First, an unknown fraction of interlopers biases the estimates of size and redshift distribution which are needed to compute the LF. The lack of spectroscopic redshift also introduces degeneracies in the SED fitting that make the masses and other physical properties poorly determined.  Finally only spectroscopic observations of high-z Ly$\\alpha$ emission have the potential to provide a definitive answer to the major question of how and when reionization occurred \\citep{Dijkstra2011,Dayal2010}. The \\lya\\ emission  line is a powerful diagnostic since  it is easily erased by neutral gas outside galaxies. Its observed strength in distant galaxies is therefore a sensitive probe  of the latest time when reionization was completed (e.g. Robertson et al. 2011). Indeed some recent studies pointed to a decrease in the number of \\lya\\ emitters (LAEs) at z$\\sim$7 \\citep{Ota2010} and failed to  detect bright LAEs at even higher redshift \\citep{Hibon2010}. In a recent work we also found a significant lack of intense \\lya\\ in ultradeep spectra of candidate  $z\\sim 7$ Lyman break galaxies \\citep{Fontana2010} (F10 from here on). All these works seem to indicate that we might be approaching the epoch when the intergalactic medium was increasingly neutral.  To address these issues we have carried out a systematic follow-up of our sample of z$\\sim 7$ candidates: galaxies were initially selected as z-dropouts in three independent fields, the GOODS-South field (Giavalisco et al. 2004), the New Technology Telescope Deep Field (NTTDF, Arnouts et al. 1999, Fontana et al. 2000) and the BDF-4 field (Lenhert \\& Bremer 2003) from deep near-IR  Hawk-I observations, covering a total area of 200 arcmin$^2$, complemented by  WFC3 observations \\citep{Castellano2010a,Castellano2010b}. We then carried out deep spectroscopic observations with FORS2. The first results were published in F10 where we reported the confirmation of a weak  Ly$\\alpha$ emission line at $z=6.97$ in the GOODS-South field, and in Vanzella et al. 2011 (V11) were we presented  two galaxies in the BDF-4 field, showing  bright Ly$\\alpha$ emission at  $z \\approx 7.1$.  In this final paper we present the results on the z-dropouts obtained in the NTT Deep Field, where we confirm two further galaxies at z$\\geq 6.7$ through the presence of \\lya\\ emission. We also present and briefly discuss the  results obtained on the i-dropout  galaxies that were observed in the three separate fields as secondary targets.  Finally we give an overview of the results obtained on the entire z-dropout sample and  compare them  to the expectations based on lower redshift observations to determine the redshift evolution of the fraction of \\lya\\ emitters in Lyman break galaxy samples. We discuss the implications of our results on the epoch of reionization.   All magnitudes are in the AB system, and we adopt $H_0=70$~km/s/Mpc, $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "We have presented and discussed spectroscopic observations of a sample of 20 z-dropout galaxies  selected from our deep and wide Hawki-I imaging program (C10a,C10b). We confirm the redshifts of 5 galaxies at $6.7 < z < 7.1$, through the presence of a \\lya\\ emission line. Only two of these galaxies have bright \\lya\\ and relatively high EWs. The number of galaxies detected in our survey is considerably smaller than what is expected from lower redshift surveys: the trend for an increasing fraction of \\lya\\ emission in LBGs that was found to hold from redshift 3 to 6 by previous studies (S11) is  halted and most probably reverted from z=6 to z=7.  If we assume that the intrinsic EW distribution of Ly$\\alpha$ is that same at $z \\sim 7$ as has been observed at $z \\sim 6$, the discrepancy between the predicted and measured EW distributions in our z-dropout sample is statistically significant. Our result would be marginally consistent with no evolution only if as many as half of the unconfirmed galaxies were interlopers. The required interloper fraction would be much larger than what we observe in our sample of i-dropouts, where interlopers are less than 18\\%, as well as what is commonly found in lower redshift samples (e.g. S10, V09).  These results extend and confirm the those of  F10, with a sample that is almost three times larger.  The  low detection rate, and the modest EW found can be explained in terms of  a rapid evolution of the neutral hydrogen fraction from z$\\sim 6$ to z$\\sim 7$. Assuming that the Universe was completely reionized at z=6, our result point are consistent with a change of  neutral hydrogen of the order of $\\Delta\\chi_{HI} \\sim 0.6$ in a relatively short time $\\Delta z \\sim 1$. Indeed, there is no evidence at all that the Universe is completely ionized at $z=6$, e.g. Mesinger (2010) argues that, since reionization is expected to be highly inhomogeneous, much of the spectra pass through just the ionized component of the intergalactic medium (IGM) even for non-negligible values of neutral hydrogen fraction (volume mean) (see also McGreer et al. 2011) However  if we want to lower substantially the $\\Delta\\chi_{HI}$ implied by our observations, then we need the Universe at $z=6$ to be  substantially neutral($>10$\\%).   In any case, even if the models uncertainties are large, the absence of strong $Ly\\alpha$ emission lines in our high  redshift galaxies  is very striking, and together with the very low number  of LAEs found in all recent deep  narrow band surveys at z$\\geq 7$ points to some drastic change that took place around that redshift. To better quantify the increase in the hydrogen neutral fraction at these epochs in the near future we will exploit the deep multi-wavelength CANDELS dataset (Grogin et al. 2011, Koekemoer et al. 2011) which will provide large numbers of candidate $z\\sim7 $ galaxies for further spectroscopic follow-up campaigns."
    },
    "1107/1107.3145_arXiv.txt": {
        "abstract": "We study relations between stellar mass, star formation and gas-phase metallicity in a sample of 177,071 unique emission line galaxies from the SDSS-DR7, as well as in a sample of 43,767 star forming galaxies at $z=0$ from the cosmological semi-analytic model \\textsc{L-Galaxies}. We demonstrate that metallicity is dependent on star formation rate at fixed mass, but that the trend is \\textit{opposite} for low and for high stellar mass galaxies. Low-mass galaxies that are actively forming stars are more metal-poor than quiescent low-mass galaxies. High-mass galaxies, on the other hand, have lower gas-phase metallicities if their star formation rates are small. Remarkably, the same trends are found for our sample of model galaxies. By examining the evolution of the stellar component, gas and metals as a function of time in these galaxies, we gain some insight into the physical processes that may be responsible for these trends. We find that massive galaxies with low gas-phase metallicities have undergone a gas-rich merger in the past, inducing a starburst which exhausted their cold gas reservoirs and shut down star formation. Thereafter, these galaxies were able to accrete metal-poor gas, but this gas remained at too low a density to form stars efficiently. This led to a gradual dilution in the gas-phase metallicities of these systems over time. These model galaxies are predicted to have lower-than-average gas-to-stellar mass ratios and higher-than-average central black hole masses. We use our observational sample to confirm that real massive galaxies with low gas-phase metallicities also have very massive black holes. We propose that accretion may therefore play a significant role in regulating the gas-phase metallicities of present-day massive galaxies. ",
        "introduction": "\\label{sec:Introduction} \\LARGE{M}\\normalsize etals are ubiquitous throughout galaxies. They are synthesised in stars and liberated into the interstellar medium (ISM) when stars shed their outer gaseous envelopes towards the end of their lives, and in some cases also into the intergalactic medium (IGM) when the highest mass stars explode as supernovae. The amount of metals in the diffuse gas around galaxies also determines the rate at which it is able to cool and form stars. Metallicity is, therefore, one of the key physical properties of galaxies, and understanding the processes that regulate the exchange of metals between stars, cold interstellar gas, and diffuse surrounding gas can help us understand the physical processes that govern galaxy evolution in general.  The metallicity of stars and gas in galaxies is known to correlate strongly with their luminosities, circular velocities and stellar masses (e.g. \\citealt{L79,G02,T04,G05}). However, the physical processes that drive these correlations are not yet fully understood.  \\citet{MB71} and \\citet{L74} first suggested that interstellar gas can be driven out of galaxies by supernova explosions as \\textit{galactic outflows}. They predicted that galaxies of smaller mass have lower metal abundances because their lower escape velocities allow freshly enriched gas to be more efficiently removed.  The input of energy from supernova explosions is now routinely incorporated into hydrodynamical simulations of galaxy formation, either in the form of thermal heating or `kinetic feedback', whereby radial momentum kicks are imparted to particles surrounding sites of star formation in the galaxy. Although these simulations are now able to demonstrate that galactic outflows can yield a good match to the observed mass-metallicity relation (e.g. \\citealt{K07,S08}), this does not mean that outflows are the only process at work in regulating the metallicities of real galaxies.  \\citet{B07} argue that a \\textit{mass-dependent star formation efficiency} (SFE) is required to explain observations, in addition to supernova feedback. In this scenario, less massive galaxies convert their gas reservoirs into stars over longer timescales than more massive galaxies. Therefore, less massive galaxies have higher gas-to-stellar mass ratios and are consequently less metal-rich. Motivated by their work on momentum-driven winds in smoothed particle hydrodynamics (SPH) simulations, \\citet{FD08} have also suggested that both metal-rich outflows at all masses and a variable star formation efficiency must play roles in explaining the observed mass-metallicity relation for nearby galaxies.  Alternatively, \\citet{D04} suggest that \\textit{metal-poor infall} can regulate metallicity in disc galaxies. Given that lower mass galaxies tend to have lower star formation rates (due to the fact that their cold gas densities are lower), a net dilution takes place when the time-scale for star formation falls below the time-scale for accretion of metal-poor gas. This would drive down the low mass end of the mass-metallicity relation.  Finally, \\textit{variations in the initial mass function} (IMF) have also been cited as another factor that might influence the mass-metallicity relation. \\citet{KWK07} propose that a SFR-dependent (and therefore stellar mass-dependent) IMF causes different galaxies to produce different effective oxygen yields. This hypothesis is based on the premise that most stars form in stellar clusters and that smaller, less actively star-forming galaxies are dominated by clusters with lower masses, containing a smaller fraction of massive, oxygen-producing stars. We note, however, that there is little observational evidence for systematic IMF variations between stellar clusters (e.g. \\citealt{El06,FDK11}).  The fact that the interpretation of the mass-metallicity relation is subject to considerable ambiguity has prompted a number of authors to consider alternative ways of quantifying metallicity in galaxies. For example, higher dimensional relations that include additional physical properties could provide better constraints on the processes that regulate metallicity. Recently, \\citet{M10} have proposed such an extension to the mass-metallicity relation, called the fundamental metallicity relation (FMR). In that work, it was shown that the gas-phase metallicities ($Z$) of both local and high-redshift galaxies were dependent on both stellar mass ($M_{*}$) and star formation rate (SFR). The FMR provides a prediction of the metallicity of local galaxies with a 1$\\sigma$ scatter of only $\\sim 0.05$ dex. This is a substantial improvement on the mean scatter of $\\sim 0.1$ dex reported by \\citet{T04} for the $M_{*}$-$Z$ relation. Metallicity was also found to be strongly dependent on SFR at low stellar masses, but only very weakly dependent on SFR at high stellar masses.  \\citet{M10} further showed that the FMR describes galaxies out to $z\\sim 2.5$. The observed evolution of the $M_{*}$-$Z$ relation was therefore attributed to migration of galaxies along the FMR plane to higher masses and lower star formation rates over time.  Prompted by these findings, in this work we study higher dimensional relations between metallicity and a variety of physical galactic properties. We examine whether the $M_{*}$-$Z$ relation exhibits additional dependences on SFR, specific SFR (sSFR) and $M_{\\textnormal{gas}}/M_{*}$. Comparisons are made between the results from observational data and the predictions from semi-analytic models of galaxy formation implemented within a high resolution simulation of the evolution of dark matter in a `concordance' $\\Lambda$CDM cosmology.  In Section \\ref{sec:The Samples} we present our observational sample, extracted from the latest Sloan Digital Sky Survey (SDSS) spectroscopic data. In Section \\ref{sec:PDMs}, we explain how we obtain estimates of SFR and $Z$ for our sample. In Section \\ref{sec:The FMR in observations} we study the dependence of the $M_{*}$-$Z$ relation on SFR, and the dependence of the SFR-$Z$ and sSFR-$Z$ relations on $M_*$. In Section \\ref{sec:The FMR in L-Galaxies} we discuss the latest version of the semi-analytic galaxy formation code \\textsc{L-Galaxies} \\citep{G10}, implemented on the \\textsc{Millennium-II} dark matter N-body simulation \\citep{BK09}, and describe how we select samples of simulated galaxies for comparison with the observational data. In Section \\ref{Model Results} we present these comparisons and show how trends in the relations between stellar mass, star formation rate and metallicity can be understood in terms of the prescriptions for gas accretion, star formation, supernova and AGN feedback that are implemented in the model. In Section \\ref{sec:Discussion}, we discuss the viability of different physical mechanisms in regulating the metallicity of galaxies. Finally, in Section \\ref{sec:Conclusions}, we summarize our results. ",
        "conclusions": "\\label{sec:Conclusions} We have shown that the gas-phase metallicities of galaxies are dependent on their star formation rate. This is also true at high masses, where highly star forming galaxies are seen to have higher metallicities -- the opposite trend to that seen at low masses. However, despite this dependence, a projection onto the $M_{*}$-SFR-$Z$ space that combines $M_{*}$ and SFR does little to reduce the scatter in $Z$ compared to the $M_{*}$-$Z$ relation.  We also demonstrate the significance of metallicity derivation methods when assessing the relation between $M_{*}$, SFR and $Z$. Strong line ratio diagnostics, such as R$_{23}$, provide significantly different metallicity estimates to Bayesian techniques which utilise emission line fluxes. These differences are greater for low star forming galaxies. Although we believe that the Bayesian technique used for our Sample T2 provides a more robust measurement of the global gas-phase metallicity in local, high-metallicity, star forming galaxies, it remains unclear to what extent this is so.  In Section \\ref{Model Results}, we show that a high-mass SFR-dependence is also present in our model sample. This is due to a turnover in the mass-metallicity relation, caused by a gradual dilution of the gas phase in some galaxies, triggered by a gas-rich merger which shuts down subsequent star formation without impeding further cooling. We have proposed that similarities between these low-sSFR model galaxies and those observed at $z=0$, such as their larger-than-average black hole masses, leave open the possibility that such a process is also driving the SFR-dependence seen at high masses in real galaxies. If this is the case, then physical processes other than mass-dependent outflows must also be playing a part in regulating metallicity. Our model indicates metal-poor galactic infall as a likely candidate."
    },
    "1107/1107.1189_arXiv.txt": {
        "abstract": "We present predictions for the radio pulses emitted by extensive air showers using ZHAireS, an AIRES-based Monte Carlo code that takes into account the full complexity of ultra-high energy cosmic-ray induced shower development in the atmosphere, and allows the calculation of the electric field in both the time and frequency domains. We do not presuppose any emission mechanism and our results are compatible with a superposition of geomagnetic and charge excess radio emission effects. We investigate the polarization of the electric field as well as the effects of the refractive index $n$ and shower geometry on the radio pulses. We show that geometry, coupled to the relativistic effects that appear when using a realistic refractive index $n>1$, play a prominent role on the radio emission of air showers. ",
        "introduction": "In the last decades, the study of ultra high energy cosmic rays (UHECR) has been one of the most active areas in astroparticle physics~\\cite{springeruhecrreview}. The detection of extensive air showers (EAS) created by UHECR in the atmosphere has been accomplished mainly with two detection methods. The first consists on detecting the particles of the cascade reaching the ground using an array of particle detectors. The second method uses telescopes to detect the fluorescence photons emitted by the particles in the shower as they travel through the atmosphere. Only the latter method is able to directly measure the EAS longitudinal development in the atmosphere, but is subject to a very low duty cycle ($ \\sim10\\%$), since it can only be used in clear moonless nights. The surface array and fluorescence techniques are used simultaneously at the Pierre Auger cosmic ray observatory~\\cite{auger}.  Radio emission from EAS was first observed by Jelley et al.~\\cite{jelley} in 1964.  Theoretical and experimental research in this area was very active in the sixties and early seventies. We refer the reader to a review by Allan~\\cite{allan} and references therein. Nevertheless, interest in this technique declined in the late seventies, mainly due to radio interference~\\cite{lopes}. Recent developments in high speed electronics and information technology have renewed interest in the detection of radio emission in the MHz range from the particles in EAS. The radio detection technique is, in principle, sensitive to the longitudinal development of the shower, like the Fluorescence technique, but its duty cycle is much higher, since it could be used anytime except during thunderstorms.  This resurgence of the radio technique has taken form as new experiments have been developed, such as CODALEMA~\\cite{codalema}, LOPES~\\cite{lopes,lopesgeo,lopesldf} and AERA~\\cite{aera}, accompanied by new calculations of the radio emission in EAS, which include analytical techniques with different levels of sophistication~\\cite{Huege_analytical,Scholten_MGMR,Montanet_radio,Ardouin_Coulomb}, Monte Carlo methods~\\cite{ReAires,Konstantinov,REAS2,REAS3,SEIFAS,zhaires_air11} and semi-analytical methods~\\cite{Scholten_MGMR_sims}.  In this work we present ZHAireS~\\cite{zhaires10,zhaires_air11} (ZHS+AIRES), a new simulation of radio emission in EAS that combines the full shower simulation capabilities of AIRES~\\cite{aires} with specific algorithms developed to calculate the electric field emitted by particles in dense media showers, implemented in the well tested ZHS code~\\cite{ZHS92,ARZ10}. These algorithms are obtained from first principles, so no emission mechanism or model is presupposed. The radio emission calculations are done in parallel to the AIRES shower simulation: As each charged particle in the shower is propagated by AIRES in steps, each propagation step is taken as a single particle track and its contribution to the radio emission is calculated and added to the total electric field in both the time and frequency domains. This procedure naturally accounts for interference effects associated to the different space-time positions of the particles in the shower.  Other Monte Carlo simulations of radio emission of air showers exist, such as REAS~\\cite{REAS2,REAS3}, where the electron and positron tracks from CORSIKA~\\cite{corsika} simulations are first histogrammed and then used to generate random $e^{\\pm}$ trajectories, which in turn are used for the radio emission calculations in the time domain only. An earlier version of this code, REAS2~\\cite{REAS2}, was based only on geosynchrotron emission, but was shown to be inconsistent with later simulations~\\cite{REAS3,zhaires_air11}, which in turn are based on a pretty generic algorithm~\\cite{allan} that has been used for a long time to simulate pulses generated in dense media~\\cite{ZHS92,ARZ10}. Another difference is that ZHAireS uses a model for the variation of the refractive index with altitude, while REAS uses a fixed refractive index equal to unity in its calculations.   This paper is organized as follows: In section \\ref{sec:emission} we give a short overview of the main radio emission mechanisms in air showers and the specific formalism used for emission calculations in ZHAireS, including a model for the variation of the refractive index with altitude. In section \\ref{sec:1Dmodel} we develop a one-dimensional toy model useful for understanding the main characteristics of the emission and stress the prominent role played by geometry. In section \\ref{sec:results} we show results of ZHAireS simulations, discussing the influence of the distance from the antenna to the shower core, the refractive index and the shower zenith angle on the electric field pulse. We also discuss the spectrum of the radio emission. In section \\ref{sec:polarization} we analyze the polarization of the electric field and discuss how it can be used to separate the contributions from  various emission mechanisms. Finally, in section \\ref{sec:conclusion} we conclude the paper.        \\section {Radio Emission in Air Showers} \\label{sec:emission} \\subsection{Main emission mechanisms} \\label{sec:mechanisms} In air showers, the dominant mechanism responsible for the radio emission is believed to be the deflection by the geomagnetic field of electrons and positrons in the shower~\\cite{codalemageo,lopesgeo,falckenature}. Several approaches have been developed to model this emission. Since the direction of the Lorentz force depends on charge, it leads to a spatial separation of electrons and positrons in the shower, that can be thought of as a moving macroscopic dipole and a transverse current traveling through the atmosphere at a speed $v \\approx c$ along with the shower front, and which has been the basis for macroscopic calculations~\\cite{kahnlerchegeo,Scholten_MGMR,Scholten_MGMR_sims}. Another approach, which is adopted in this work, is to calculate the emission for each particle trajectory, i.e. a microscopic approach~\\cite{ZHS92,ARZ10,zhaires10,zhaires_air11,REAS3}.  The geomagnetic mechanism has a quite clear signature. Its polarization is anti-parallel to the direction of the Lorentz force, i.e. in the direction of $-\\vec{\\beta}\\times\\vec{B}$, where $\\vec{\\beta}$ is the speed of the particle in $c$ units and $\\vec{B}$ the geomagnetic field~\\cite{kahnlerchegeo,lerchenature,codalemavxb,scholtenarena2010}.  Another emission mechanism also thought to be important in EAS is the Askaryan effect. It was first proposed by Askaryan~\\cite{Askaryan62}, who suggested that coherent radiation could be emitted by showers in which a charge excess develops. This mechanism, which dominates the radio emission of showers in dense media, is also known as the charge excess mechanism. For a particle following a straight track at constant speed, it can be shown \\cite{Jackson} that: \\begin{equation} \\vec{E}_{rad}\\propto -e [ \\hat{u} \\times ( \\hat{u} \\times \\vec{\\beta} ) ] \\end{equation}  \\noindent So the electric field emitted by  particles traveling along the shower axis lies in the plane defined by the direction of the shower axis ($\\vec{\\beta}$) and the observation direction $\\hat{u}$. Furthermore, it is perpendicular to $\\hat{u}$ and its direction depends on the charge of the particle, pointing towards (away from) the shower axis for negative (positive) charges \\cite{ZHS92,ARZ10}. This means that if there is no charge excess in the shower, the net electric field is zero. However it is well known that knock-on interactions and Compton scattering incorporate electrons from molecules of the medium into the shower, leading to an excess of moving negative charges \\cite{Askaryan62}, which in turn are responsible for a net electric field with a radial polarization w.r.t. the shower core.  \\subsection{Radio emission calculation in the \\ZHAireS code} \\label{newcode}  The algorithms used for the calculation of radio emission in ZHAireS are based on ZHS algorithms~\\cite{ZHS92,ARZ10}. These were derived from first principles, namely from the Lienard-Wiechert potentials, and thus do not assume any emission mechanism. The full derivation can be seen in~\\cite{ZHS92} for frequency domain calculations, and~\\cite{ARZ10} for the time domain. The trajectories of shower particles are divided in tracks, which can be made arbitrarily small. Both the speed of the particle and the direction to the observer are assumed not to vary over the track. In Fig. \\ref{fig:singletrack} we show a schematic picture of such a track. A convenient division of trajectories in small tracks is already performed by Monte Carlo simulations of shower development in order to propagate the particles in the shower. The ZHS algorithm is then used to calculate the contribution of each track to the net emission of the shower, accounting for any interference between tracks. This approach automatically takes into account the contribution to the electric field due to the start, end, and any change in the direction and energy of each particle track. Thus any kind of deflection, scattering, creation or annihilation of a charged particle, due to any physical process used to simulate the shower is taken into account in the radio emission\\cite{zhaires_air11,zhaires10,endpoint}.    \\begin{figure} \\begin{center} \\scalebox{0.8}{ \\includegraphics{singletrack.eps}} \\caption{Scheme of a single particle track. Speed and direction to the observer are assumed constant inside each track.} \\label{fig:singletrack} \\end{center} \\end{figure}  In the original ZHS algorithms~\\cite{ZHS92,ARZ10}, the field calculation is performed in the Fraunhofer approximation. In this case the observer ``sees'' the radiation arriving from all the points in the shower at the same angle, and only changes in the phase of the contribution from each point are taken into account through a simple projection. This approach works well for showers in dense media, such as ice, since the size of the shower is much smaller than the typical distance to the antennas. But in the case of air showers, the distance to the observer is usually of the order of the shower size, and thus the Fraunhofer approximation breaks down. In order to make the method valid for the closer observers in air showers, we allow the distance and direction to the observer to change from one track to the next. The direction and distance to the antenna are still taken to be constant inside a single track (see Fig.~\\ref{fig:singletrack}), i.e. we still use the Fraunhofer approximation inside each track. It can be shown~\\cite{inpreparation} that this procedure reproduces the expected behavior $E\\propto 1/\\sqrt{R}$ of the field~\\cite{hunearfield}, where $R$ is the distance between the charge and the observer. Since we still use the Fraunhofer approximation inside a single track of length L, the following condition has to be satisfied:   \\begin{equation} \\frac{L^2\\sin^2{\\theta}}{R}<\\frac{\\lambda}{2\\pi}\\;\\; \\label{eq:fraunhoffer} \\end{equation}   \\noindent where $\\lambda$ is the wavelength of the emission. In the simulation, the vast majority of tracks satisfy this condition for frequencies up to 300 MHz, which is much higher than the frequency of the maximum of the emission, $\\sim 1-30$ MHz. If a particular track in the simulation does not satisfy this condition, e.g. a track passes very close to an antenna, it is further divided into several sub-tracks, and the field calculation is performed using a different R and $\\theta$ for each sub-track.   The positions $\\vec{x}_{1,2}$, times $t_{1,2}$ and kinetic energies $E_{1,2}$ for the beginning and end points of the track are obtained directly from AIRES. This, along with the position $\\vec{x}_{ant}$ of the antenna is all that is needed for the calculation of the contribution to the vector potential $\\vec{A}(t,\\hat{u})$ due to a given track: \\begin{equation} \\vec{A}(t,\\hat{u})=\\frac{\\mu e}{4\\pi R c}\\vec{\\beta}_\\perp\\frac{\\Theta(t-t^{det}_1)-\\Theta(t-t^{det}_2)}{1-n\\vec{\\beta}\\cdot\\hat{u}} \\label{eq:timefresnel} \\end{equation} \\noindent where $\\vec{u}=R\\hat{u}$  is the vector from the middle point of the track to the antenna, $\\vec{\\beta}=\\vec{v}/c$, $\\vec{\\beta}_\\perp=-\\left[\\hat{u}\\times(\\hat{u}\\times \\vec{\\beta}) \\right]$ is the projection of $\\vec{\\beta}$ onto a plane perpendicular to $\\hat{u}$, $t^{det}_{1,2}=t_{1,2}+nR/c-n\\vec{\\beta}\\cdot\\hat{u}~(t_{1,2}-t_0)$ are the retarded (detection) times for the beginning and end of the track, respectively, $t_0=(t_1+t_2)/2$ is the average time for the track and $\\Theta(x)$ is the Heaviside step function\\footnote{In the formalism, as described in ~\\cite{ARZ10}, the limit of eq. (\\ref{eq:timefresnel}) for $(1-n\\vec{\\beta}\\cdot\\hat{u})\\rightarrow 0$ is used  instead of eq. (\\ref{eq:timefresnel}) if the denominator vanishes.}. Note that Eq.~(\\ref{eq:timefresnel}) is written in the radiation gauge, since we disregard the static term of the Lienard-Wiechert potentials in its derivation~\\cite{ARZ10}.  To obtain the radio signal at each antenna for the shower as a whole, the contributions of each particle track to the vector potential $\\vec{A}(t)$ are added up\\footnote{Note that interference effects are taken into account automatically, since the contribution of each track to the vector potential is related to a specific retarded time.}. This net vector potential is then differentiated with respect to time to obtain the net electric field as a function of time for each antenna. In the far field, this formalism is equivalent to the one used in~\\cite{REAS3}, as shown in~\\cite{endpoint}.  Besides the field calculations in the time domain, ZHAireS can calculate the radio emission in the frequency domain as well. The radiation term of the electric field is $\\vec{E}(t)=-\\partial \\vec{A}/ \\partial t$. Applying this to the vector potential expression (eq. \\ref{eq:timefresnel}) we obtain~\\cite{ARZ10}: \\begin{equation} \\vec{E}(t,\\hat{u})=-\\frac{\\mu e}{4\\pi R c}\\vec{\\beta}_\\perp\\frac{\\delta(t-t^{det}_1)-\\delta(t-t^{det}_2)}{1-n\\vec{\\beta}\\cdot\\hat{u}} \\label{eq:timefresnel-EF} \\end{equation}  In the ZHS formalism~\\cite{ZHS92,ARZ10}, we use the following convention for the Fourier transform: \\begin{equation} \\tilde{f}(\\omega)=2\\int_{-\\infty}^{\\infty}{f(t) e^{i\\omega t}dt} \\label{eq:ft} \\end{equation} Applying this Fourier transform convention to Eq.~(\\ref{eq:timefresnel-EF}) we obtain~\\cite{ZHS92}: \\begin{equation} \\vec{E}(\\omega,\\hat{u})=-\\frac{\\mu e}{2\\pi R c}\\vec{\\beta}_\\perp\\frac{e^{i\\omega(t-t^{det}_1)}-e^{i\\omega(t-t^{det}_2)}}{1-n\\vec{\\beta}\\cdot\\hat{u}} \\label{eq:freqfresnel} \\end{equation} \\noindent which is the expression we use for the Fresnel regime frequency-domain calculations in ZHAireS. The contributions of each track to $\\vec{E}(\\omega)$ are added up to obtain the net spectrum at each antenna.  \\subsection{Variable refractive index} \\label{sec:n(h)}  As stated before, in ZHAireS we use a model for the variation of the refractive index $n$ with altitude in the atmosphere. In fact there is a strong dependence of $n$ on temperature, pressure and humidity. To take the dependence with altitude into account we use an exponential model for the variation of the refractivity ${\\cal R}$ with height $h$, motivated by the exponential decrease of density with altitude, \\begin{equation} {\\cal R}(h)={\\cal R}_s \\exp{(-K_r h)} \\label{eq:n(h)} \\end{equation} where ${\\cal R}(h)=[n(h)-1]\\times 10^6$ is the refractivity at an altitude $h$ in km and we used ${\\cal R}_s={\\cal R}(h=0)=325$ and $K_r=0.1218~{\\rm km^{-1}}$. These values reproduce the refractivity calculated in~\\cite{Gerson} up to $h\\sim10$ km within less than $\\sim 1\\%$. The values in \\cite{Gerson} take into account the humidity dependence of the refractivity,  which increases the refractivity at low altitudes where the effect of a variable index of refraction is most important, since the shower has a larger number of particles. At higher altitudes, the exponential model slightly overestimates the refractivity, but the effect of the variable index of refraction at these altitudes is much less important, since above 20 km the shower has barely started developing and the number of particles is small. It is also worth noting that since ${\\cal R}$ depends on temperature and humidity, there are large seasonal and even daily variations in the refractivity that are associated to atmospheric conditions and cannot thus be accurately described by any model\\footnote{If the humidity as a function of height can be monitored using e.g. a LIDAR, then more accurate values of the refractivity for specific atmospheric conditions could be calculated by adding a humidity term to Eq. (\\ref{eq:n(h)}), as described in~\\cite{Gerson}.}.  The variable atmospheric refractive index would in principle make the radio emission follow a curved path. However, in \\cite{Scholten_MGMR_sims} it was shown that the effect of the deviation from a straight path on the time structure of the pulse is negligible, and thus in ZHAireS we assume that the radio emission follows a straight path from the emission point to the antenna. However, we explicitely take into account the effect of a variable refractive index in the propagation time of the signal. For this purpose we calculate an effective refractive index $n_{eff}$ for each particle track in the ZHAireS shower simulation: \\begin{equation} n_{eff}=1+{\\cal R}_{eff}\\times10^{-6},\\;\\;\\;{\\cal R}_{eff}=\\frac{1}{R}\\int_{0}^{R}{{\\cal R}(h)~dl} \\label{eq:neff} \\end{equation} \\noindent where $R$ is the distance from the track to the observer, and $dl$ is an infinitesimal length along that path whose altitude $h$ varies along the path. It is important to note that while $n_{eff}$ is used for the calculation of the retarded times, $n(h)$ is used for the angular dependence of the emission, given by the denominator of  Eqs. (\\ref{eq:timefresnel}) and (\\ref{eq:freqfresnel}), i.e. the Cherenkov angle depends only on the refractive index at the emission altitude.   We expect the effect of the variable refractive index to be more important at relatively low altitudes, where $n$ is larger. Since $n$ in our model varies between $n\\sim 1$ at high altitudes and $n\\sim1.000325$ close to ground, we expect the effect of the variable $n$ on the pulse characteristics discussed in section \\ref{sec:1Dmodel} (such as start-time, duration and compression in time) to be between the effects obtained for $n=1$ and $n=1.0003$. This can be clearly seen in Figs.~\\ref{fig:tdetxdepth} and \\ref{fig:fc-varn} below. ",
        "conclusions": "\\label{sec:conclusion}  In this work we present predictions for the radio pulse emitted by extensive air showers. Our results are obtained using the ZHAireS Monte Carlo, an AIRES-based code that takes into account the full complexity of ultra-high energy cosmic-ray induced shower development in the atmosphere, and allows the calculation of the electric field in both the time and frequency domains based on the algorithms developed in \\cite{ZHS92,ARZ10}. Although our approach does not presuppose any a priori emission mechanism, our results confirm that the emission at radio frequencies can be understood as the superposition of radiation from the charge separation induced by the magnetic field of the Earth (geomagnetic effect), and that coming from the net excess negative charge evolving as the shower develops in the atmosphere (Askaryan effect).  We have pointed out the relevance of the refractive index in the time structure and intensity of the radio pulses, especially at short distances to the shower axis. Another interesting work~\\cite{scholtenprl}, developed independently and in parallel with ours, deals with similar issues. The refractive index determines the angular distribution of the radiation at the emission point, as well as its propagation through the atmosphere and non trivial relativistic effects arise due to the refractive index $n>1$. We have developed a simple 1-dimensional model to address the characteristics of the radio pulse, which qualitatively allows us to interpret the pulse height, width and its dependence on the refractive index and the distance from antenna to shower core \\cite{scholtenprl}. The intensity of the radio pulses is typically highest for those observers that see a region of the shower containing a large number of particles (such as shower maximum) with an angle close to the Cherenkov angle. There is a non-trivial interplay between the distance from the emission point in the shower to the observation point, the angle between the particle direction and the observer, and the number of particles in the region of the emission, whose effects can only be accurately determined with Monte Carlo simulations such as those developed in this work. These key elements determine for instance that far from the shower axis inclined showers typically induce a more intense radiation than vertical ones \\cite{Gousset_inclined} as we show with our ZHAireS simulations.  We have shown that the frequency at which the emission spectrum is maximum, which is in the range $\\sim3$-$30$~MHz for $0<r<400$ m, decreases as the distance $r$ from the antenna to the shower core increases (for $r>100$~m see also \\cite{Scholten_MGMR,REAS3,compREAS3MGMR}). So most experiments, which typically are only sensitive to frequencies above $30$~MHz, in fact only measure the incoherent part of the spectrum, even very close to the shower core. We also observed that the signal at higher frequencies decreases more rapidly with $r$ than at the fully coherent 1 MHz range, where we observed a maximum in the emission at around $r\\sim 100$ m, similar to the one reported in \\cite{scholtenprl}. Our results also suggest that this maximum is dependant on frequency and on observer direction w.r.t. the shower core. We have also explored the polarization properties of the radiation, confirming the expectation that the radiation is mainly polarized in the opposite direction to the Lorentz force induced by the magnetic field of the Earth~\\cite{kahnlerchegeo,lerchenature,codalemavxb,scholtenarena2010}. Far from the shower core, our simulations show that the polarization is compatible with the presence of a significant amount of radiation due to the Askaryan effect, as in~\\cite{REAS3}. We have also shown, using a very simple model, that as the Askaryan component of the emission becomes dominant over the geomagnetic one at larger distances to the core, the parameter ${\\cal R}_p$ changes its periodicity on the azimuthal angle of the observer. When both, the geomagnetic and Askaryan components are significant, the periodicity is $360^\\circ$, but when the emission is dominated by the Askaryan effect, a $180^\\circ$ periodicity can be clearly appreciated. In inclined showers we have stressed the role played by the relative orientation of the shower axis and the magnetic field on the polarization. A significant vertical component of the electric field arises in inclined showers which calls for detection instruments able to observe it."
    },
    "1107/1107.0977_arXiv.txt": {
        "abstract": "The relativistic wind of obliquely-rotating pulsars consists of toroidal stripes of opposite magnetic field polarity, separated by current sheets of hot plasma. By means of two- and three-dimensional particle-in-cell simulations, we investigate particle acceleration and magnetic field dissipation at the termination shock of a relativistic striped wind. At the shock, the flow compresses and the alternating fields annihilate by driven magnetic reconnection. Irrespective of the stripe wavelength $\\lambda$ or the wind magnetization $\\sigma$ (in the regime $\\sigma\\gg1$ of magnetically-dominated flows), shock-driven reconnection transfers all the magnetic energy of alternating fields to the particles, whose average Lorentz factor increases by a factor of $\\sigma$ with respect to the pre-shock value. The shape of the post-shock spectrum depends primarily on the ratio $\\lambda/(r_L\\sigma)$, where $r_L$ is the relativistic Larmor radius in the wind. The spectrum becomes broader as the value of $\\lambda/(r_L\\sigma)$ increases, passing from a relativistic Maxwellian to a flat power-law tail with slope around $-1.5$, populated by particles accelerated by the reconnection electric field. Close to the equatorial plane of the wind, where the stripes are symmetric, the highest energy particles resulting from magnetic reconnection can escape ahead of the shock, and be injected into a Fermi-like acceleration process. In the post-shock spectrum, they populate a power-law tail with slope around $-2.5$, that extends beyond the flat component produced by reconnection. Our study suggests that the spectral break between the radio and the optical band in Pulsar Wind Nebulae can be a natural consequence of particle acceleration at the termination shock of striped pulsar winds. ",
        "introduction": "\\label{sec:intro} The broadband spectrum of Pulsar Wind Nebulae (PWNe), extending from the radio up to the gamma-ray band, is usually modeled as synchrotron and inverse Compton radiation from a nonthermal population of electron-positron pairs. It is assumed that the emitting particles are accelerated to a power-law distribution at the so-called termination shock, where the momentum flux of the relativistic pulsar wind is balanced by the confining pressure of the nebula.  However, it is still unclear how the wind termination shock can accelerate electrons and positrons (hereafter, simply ``electrons'') to the relativistic energies required to power the nebular emission.  A flat radio spectrum is the common observational feature of PWNe, whose spectral flux in the radio band scales as $F_{\\nu}\\propto\\nu^{-\\alpha_r}$, with $\\alpha_r=0.0-0.3$. For the Crab Nebula, the best studied example of a PWN, the synchrotron cooling time of radio-emitting electrons largely exceeds the age of the nebula, so one cannot exclude that such particles may be ``primordial,'' i.e., accelerated at a very early stage of the PWN evolution  \\citep{atoyan_99}. However, observations of wisps in the radio band \\citep[][]{bietenholz_01} seem to suggest that radio-emitting electrons are being accelerated now, in the same region as those responsible for the optical and X-ray emission. In this case, the observed radio spectral index \\citep[$\\alpha_r\\sim0.25$ for the Crab;][]{bietenholz_97} would imply a distribution of emitting electrons of the form $dN/dE\\propto E^{-p}$, with slope $p=2\\alpha_r+1\\sim1.5$. To explain the radio through optical emission of the Crab, the power law of shock-accelerated electrons should span at least three decades in particle energy \\citep[e.g.,][]{lyubarsky_03}. Flat electron spectra with slope $p<2$ below GeV energies are also required to model the radio emission of hotspots in radio galaxies \\citep[][]{stawarz_07} and X-ray observations of luminous blazar sources \\citep[][]{sikora_09}.  For a power-law particle spectrum with $1<p<2$, most of the particles lie at the low-energy end of the distribution, but the plasma energy content is dominated by particles at the high-energy cutoff. This  contrasts with the energy spectrum expected from Fermi acceleration in relativistic shocks, which normally yields slopes $p>2$ \\citep[e.g.,][]{achterberg_01, keshet_waxman_05}, in which case low-energy particles dominate both by number and by energy. To produce broad spectra with $p<2$, a mechanism is required that raises the mean electron energy, but leaves most of the particles at relatively low energies. In the context of pulsar winds, \\citet{hoshino_92} and \\citet{amato_arons_06} have proposed that, in a wind loaded with ions, resonant absorption of ion cyclotron waves by electrons and positrons could generate flat distributions downstream from the termination shock. However, the required ion injection rate largely exceeds what is expected from standard models of pulsar magnetospheres \\citep[e.g.,][]{arons_07}.  An alternative possibility was discussed by \\citet{lyubarsky_03}, under the assumption that the flow upstream of the termination shock consists of alternating stripes of opposite magnetic polarity, separated by current sheets of hot plasma (from now on, a ``striped wind''). For obliquely-rotating pulsars, this is the configuration expected around the equatorial plane of the wind, where the sign of the toroidal field alternates with the pulsar period.  If the wind remains dominated by Poynting flux till the termination shock \\citep[which is still an open question, see][]{lyubarsky_kirk_01, kirk_sk_03}, shock-compression of the stripes may drive annihilation of the alternating fields, a process  known as ``driven magnetic reconnection.'' The resulting transfer of energy from the fields to the particles could potentially generate flat electron distributions, with slope $1<p<2$.  The physics of magnetic reconnection can be captured self-consistently only by means of multi-dimensional particle-in-cell (PIC) simulations. Fully-kinetic PIC simulations provide a powerful tool to explore the microphysics of collisionless plasmas, since they can capture from first principles the fundamental interplay between charged particles and electromagnetic fields. In the context of relativistic magnetic reconnection in pair plasmas, most studies have explored the so-called ``undriven reconnection'' process \\citep{zenitani_01,zenitani_05, zenitani_07,zenitani_08, jaroschek_04}, where field annihilation is initiated by a transient seed perturbation to an otherwise stable current sheet. As discussed above, this is not the setup expected at the wind termination shock. Here, it is the shock-compression of the flow that steadily drives regions of opposite magnetic field polarity toward each other, causing reconnection.  Shock-driven reconnection in magnetically-dominated flows has been studied by \\citet{petri_lyubarsky_07} with analytical methods and one-dimensional (1D) PIC simulations.\\footnote{A similar 1D study, but for winds dominated by kinetic (rather than Poynting) flux, was performed by \\citet{nagata_08}.} They find a condition on the wind magnetization and the stripe wavelength needed for full dissipation of the alternating fields at the termination shock (roughly speaking, the post-shock Larmor radius should be larger than the separation between current sheets). However, in 1D all particles gain energy at the same rate, which results in a Maxwellian distribution downstream from the shock. The generation of nonthermal tails is possible only if different particles experience different energy gains, which requires multi-dimensional simulations. As a preliminary step toward multi-dimensional studies of the termination shock in a striped wind, \\citet{lyubarsky_liverts_08} have performed 2D simulations of driven magnetic reconnection. Two stripes of opposite magnetic polarity were compressed  by means of an external force, which would imitate the effect of a shock. They found that driven magnetic reconnection can produce flat nonthermal tails, with $p\\sim1$. Yet, their study was limited to a single current sheet, and most importantly it did not self-consistently account for the presence of a shock.  In this work, we explore via multi-dimensional PIC simulations the acceleration of particles at the termination shock of a striped relativistic electron-positron wind. The physics of a shock interacting with multiple current sheets is captured self-consistently in 2D and 3D, and the transfer of energy from the field to the particles via shock-driven reconnection is investigated from first principles. We find that the alternating fields are completely dissipated upon compression by the shock, and their energy is transferred to the particles, regardless of the properties of the flow. This contrasts with the conclusions of \\citet{petri_lyubarsky_07}, who found efficient field dissipation only for high magnetizations and short stripe wavelengths. We show that the reduced dimensionality of their 1D model could not correctly capture the growth of the so-called ``tearing-mode instability,'' which is of paramount importance for field annihilation in 2D and 3D configurations. We find that broad particle spectra with flat slopes ($1<p<2$) are a common by-product of shock-driven reconnection, but the extent of the power-law tail depends on the wind magnetization and the stripe wavelength.  This work is organized as follows. In \\S \\ref{sec:setup}, we discuss the setup of our simulations and the magnetic field geometry. In \\S \\ref{sec:struct}, the properties of the termination shock are investigated for one representative choice of wind parameters. We discuss both the shock structure, with focus on the physics of shock-driven magnetic reconnection, and the process of particle acceleration to nonthermal energies. In \\S \\ref{sec:cond}, we study how the particle spectrum downstream from the shock depends on the physical conditions in the wind, like magnetization and stripe wavelength. Although we mostly employ 2D simulations, in \\S\\ref{app:3d} we compare 2D and 3D results, finding good agreement. We summarize our findings in \\S \\ref{sec:disc} and comment on the application of our results to astrophysical scenarios. ",
        "conclusions": "\\label{sec:disc} We have explored by means of 2D and 3D PIC simulations the internal structure and acceleration properties of relativistic shocks that propagate in an electron-positron striped wind, i.e., a flow consisting of stripes of alternating field polarity separated by current sheets of hot plasma. This work extends the study of \\citet{sironi_spitkovsky_09}, that investigated relativistic shocks in a pair plasma carrying uniform magnetic fields.  We find that a fast MHD shock propagates into the pre-shock striped flow, compressing the incoming current sheets and initiating the process of \\textit{driven magnetic reconnection}, via the tearing-mode instability. Reconnection islands seeded by the passage of the fast shock grow and coalesce, while magnetic energy is dissipated at X-points located in between each pair of islands. When reconnection islands grow so big to occupy the entire region between neighboring current sheets, the striped structure of the flow is erased, and a hydrodynamic shock forms. Downstream from the shock, the average particle energy is larger than in the pre-shock flow by a factor of $\\simeq\\sigma$, where $\\sigma$ is the wind magnetization. In other words, the energy stored in the alternating fields has been entirely transferred to the particles via shock-driven magnetic reconnection, and the downstream fluid behaves like an unmagnetized plasma. The only field component surviving in the downstream region comes from shock-compression of the stripe-averaged pre-shock field (i.e., the non-alternating component), if non-zero.  In our 2D and 3D simulations, we find that complete annihilation of the alternating fields occurs irrespective of the stripe wavelength $\\lambda$ or the magnetization $\\sigma$, in contrast with the results of the 1D analysis by \\citet{petri_lyubarsky_07}. In Appendix \\ref{app:my} we show that multi-dimensional studies are of paramount importance to correctly capture the physics of shock-driven reconnection.  The main agent of particle energization is the reconnection electric field, as particles drift from a given X-point into the closest island. Whether all particles have comparable energy gains, or only a few particles are accelerated to high energies, and the majority stay cold, depends sensitively on the properties of the wind. We have explored the dependence of the downstream particle spectrum on the stripe wavelength $\\lambda$, the wind magnetization $\\sigma$ (or equivalently, the field strength $B_0$), and the stripe-averaged magnetic field $\\langle B_y\\rangle_\\lambda$ (or equivalently, the parameter $\\alpha$, with $\\alpha=0$ for $\\langle B_y\\rangle_\\lambda=0$, and $|\\alpha|\\rightarrow1$ in the limit $|\\langle B_y\\rangle_\\lambda|\\rightarrow B_0$). For fixed $\\alpha=0.1$, we find that the shape of the post-shock spectrum depends primarily on the combination $\\lambda/(r_L\\sigma)$, where $r_L$ is the relativistic Larmor radius in the striped wind. This is just the stripe wavelength measured in units of the post-shock plasma skin depth. For small values of $\\lambda/(r_L\\sigma)$ ($\\lesssim$ a few tens), the spectrum resembles a Maxwellian distribution with mean energy $\\langle\\gamma\\rangle\\simeq\\gamma_0\\sigma$, where $\\gamma_0$ is the bulk Lorentz factor of the pre-shock flow. In the limit of very large values of $\\lambda/(r_L\\sigma)$ ($\\gtrsim$ a few hundreds), the spectrum approaches a broad power-law tail, with a flat spectral slope $p\\simeq1.5$. In this regime, the tail extends from  $\\gamma_{\\rm min}\\simeq\\gamma_0$ up to $\\gamma_{\\rm max}\\simeq \\gamma_0 \\sigma^{1/(2-p)}$. In terms of the flow structure, the former case (i.e., a Maxwellian-like spectrum) is realized when all particles pass in the vicinity of an X-point, thus gaining energy from field dissipation. In contrast, the latter case occurs when most of the particles remain far from X-points, thus staying cold, and the energy of annihilating fields is transferred to only a few particles, which end up dominating the energy content of the flow.  As $|\\alpha|$ increases from $0.1$ up to $1$, the fraction of upstream Poynting flux in the form of alternating fields becomes smaller. As the amount of field energy available for dissipation decreases, the average particle Lorentz factor downstream from the shock recedes from $\\langle\\gamma\\rangle\\simeq\\gamma_0\\sigma$, as expected for complete field annihilation, down to $\\langle\\gamma\\rangle\\simeq\\gamma_0$, the result of  an unstriped wind. In the opposite limit $|\\alpha|\\lesssim0.1$, i.e., in the case of nearly symmetric stripes, the particles accelerated by the reconnection electric field to the highest energies can escape ahead of the shock, unaffected by the stripe-averaged field that would tend to advect them back downstream. Their counter-streaming with respect to the incoming flow can seed the filamentation (or Weibel) instability, as it happens for shocks propagating in unmagnetized plasmas \\citep{spitkovsky_05,spitkovsky_08}. In turn, the turbulence generated by the Weibel instability can mediate particle acceleration to even higher energies, via a Fermi-like diffusive process. The nonthermal tail produced by the Fermi mechanism is steeper than that given by magnetic reconnection, with characteristic slopes $p\\gtrsim2.5$. Typically, a few percent of particles are injected into the Fermi process, and the injection efficiency scales in inverse proportion to the ratio $\\lambda/(r_L\\sigma)$. With time, the nonthermal tail of Fermi-accelerated particles extends to higher energies. For $|\\alpha|\\gtrsim0.1$, the propagation of particles far ahead of the shock is inhibited by the stripe-averaged field, and Fermi acceleration is suppressed. In this case, a few percent of particles can get accelerated by the stripe-averaged electric field $\\langle E_z\\rangle_\\lambda$ while gyrating around the shock front, a process known as shock-drift acceleration \\citep[i.e.,][]{begelman_kirk_90}. The tail resulting from shock-drift acceleration has a limited energy extent, and it does not evolve in time.  We have confirmed our findings with an extensive set of convergence tests. Most importantly, as we have shown in \\S\\ref{app:3d}, 2D simulations with in-plane magnetic fields -- the setup that we have primarily adopted in this work -- can reproduce very accurately the shock physics and particle spectrum observed in 3D simulations. In other words, the current sheet folding caused by the drift-kink instability in the plane orthogonal to the field \\citep{zenitani_07} does not represent a serious constraint for the maximum particle energy achieved in the course of the reconnection process. Rather, the maximum energy of accelerated particles is limited by the time to drift from a given X-point into the closest island.   These findings can place important constraints on current models of Pulsar Wind Nebulae (PWNe), bubbles of synchrotron-emitting plasma powered by the relativistic wind of young pulsars. If the magnetic and rotation axes of the pulsar are not aligned, the wind consists of stripes of opposite magnetic field polarity, alternating with the pulsar period. The wind wavelength will be $\\lambda=2\\pi R_{LC}$, where $R_{LC}=c/\\Omega$ is the so-called light-cylinder radius, and $\\Omega$ is the angular frequency of the pulsar. The stripe-averaged field vanishes along the equatorial plane (i.e., $\\alpha=0$ there), and the striped structure persists up to a latitude equal to the inclination angle between the pulsar's magnetic and rotational axes. At higher latitudes, the wind carries a uniform (i.e., non-alternating) field, so the physics discussed in the present work does not apply \\citep[see, instead,][]{gallant_92, hoshino_92, amato_arons_06, sironi_spitkovsky_09, sironi_spitkovsky_11a}.  Our results suggest that efficient dissipation of the alternating fields should occur at the termination shock of pulsar winds, irrespective of the properties of the flow. It follows that current observational constraints on the upstream magnetization parameter \\citep{rees_gunn_74, kennel_coroniti_84, delzanna_04} should be attributed not to the total Poynting flux, but to the Poynting flux associated with the stripe-averaged magnetic field, the only component surviving downstream from the shock. The upstream flow may be Poynting-dominated, provided that most of the Poynting flux is in the form of an alternating field, which is annihilated at the shock. Thus, our results provide a solution to the so-called ``sigma problem,'' i.e., that pulsar winds should be magnetically-dominated at the light-cylinder radius, but kinetically-dominated downstream from the termination shock. The idea of shock-driven reconnection as a solution to the sigma problem was originally proposed by \\citet{lyubarsky_03}. Our work demonstrates that complete annihilation of the alternating fields occurs regardless of the properties of the wind, in the regime of relativistic magnetically-dominated flows.  In addition, the particle energy spectra extracted from our simulations can be used directly to interpret the radiative signature of PWNe. The radio spectrum of the Crab Nebula, the prototype of the class of PWNe, requires a population of nonthermal particles with a flat spectral slope ($p\\simeq1.5$), extending at least across three decades in energy, from $10^2\\unit{MeV}$ up to $10^5\\unit{MeV}$. The particle spectrum should be steeper at higher energies, with $p\\simeq2.5$, to explain the optical and X-ray flux.  Our results suggest that, for a slope $p\\simeq1.5$, the width of the electron power-law tail resulting from reconnection will be $\\simeq\\sigma^2$. To produce a particle distribution extending over three decades in energy, the wind magnetization at the termination shock needs to be $\\sigma\\gtrsim30$. Most importantly, broad particle spectra are produced by shock-driven reconnection only if $\\lambda/(r_L \\sigma)\\gtrsim$ a few tens. For pulsar winds, the particle number density scales with distance from the pulsar as $n=n_{LC}(R_{LC}/R)^2$, where  the density at the light-cylinder radius is usually written as $n_{LC}=\\kappa\\,n_{GJ}$. Here, $n_{GJ}=\\Omega B_{LC}/2\\pi ec$ is the Goldreich-Julian density \\citep{goldreich_julian_69}, and $\\kappa$ is the so-called multiplicity. Conservation of energy along the flow streamlines implies that at the termination shock $\\gamma_0 (1+\\sigma) \\kappa=\\omega_{LC}/2\\,\\Omega$, where $\\omega_{LC}=eB_{LC}/mc$. It follows that at the termination shock ($R=R_{TS}$) \\be\\label{eq:wind} \\frac{\\lambda}{r_L \\sigma}\\simeq4\\pi\\kappa \\frac{R_{LC}}{R_{TS}}~~. \\ee The constraint $\\lambda/(r_L \\sigma)\\gtrsim$ a few tens required to produce broad particle spectra is extremely challenging for current models of PWNe and pulsar magnetospheres. The radius of the Crab termination shock in the equatorial plane is inferred from X-ray observations to be $R_{TS}\\simeq0.1\\unit{pc}\\simeq5\\times10^8R_{LC}$ \\citep{hester_02}. Most available models estimate $\\kappa\\simeq10^4-10^6$ \\citep{bucciantini_11}, depending on whether or not radio-emitting electrons are included in the analysis. Based on our findings,  the resulting value of $\\lambda/(r _L \\sigma)\\lesssim0.01$ would yield a Maxwellian-like spectrum, at odds with the wide flat spectrum required by observations.  If radio-emitting electrons are produced in the equatorial plane by shock-driven reconnection, a revision of the existing estimates of $\\kappa$ is required. In this respect, we point out that the values of $\\kappa$ quoted above are averages over latitude, and one cannot exclude that particle injection into the pulsar wind is highly anisotropic, with multiplicity as large as $10^8$ along the equatorial plane. Of course, any given choice of $\\kappa$ will constrain $\\gamma_0$ and $\\sigma$ via  $\\gamma_0 (1+\\sigma) \\kappa=\\omega_{LC}/2\\,\\Omega$, a relation that holds at each latitude. From the equatorial plane, radio-emitting electrons will quickly fill the whole nebula, transported by the strong fluid motions observed in MHD models of PWNe downstream from the termination shock \\citep[e.g.,][]{komissarov_04, delzanna_04, camus_09}. This would provide a natural explanation for the lack of gradients in the radio spectral slope of the Crab \\citep{bietenholz_kronberg_92}. Alternatively, the radio part of the spectrum may be produced at higher latitudes, where the termination shock is closer to the pulsar, and the ratio in \\eq{wind} becomes larger. However, in this case one should also account for the fact that, at high latitudes, the fraction of upstream Poynting flux available for dissipation is lower (since $|\\alpha|$ is larger), which results in a narrower spectrum, everything else being fixed.\\footnote{Also, the termination shock at high latitudes is oblique with respect to the wind velocity, a configuration not studied here.}  On the other hand, the small value of $\\lambda/(r_L \\sigma)$ expected in the equatorial plane on the basis of current models of the Crab Nebula is the most favorable setup for particle acceleration via the Fermi diffusive process, with efficiency of a few percent by number. In this respect, one could interpret the optical and X-ray signature of the Crab, which requires a particle spectrum with $p\\simeq2.5$, as synchrotron emission from such Fermi-accelerated particles. Based on our results, optical and X-ray emitting electrons should be produced close to the equatorial plane of the wind, where $|\\alpha|\\lesssim0.1$. If the same equatorial wedge is responsible for accelerating the radio-emitting electrons, our findings predict that they should outnumber the optical and X-ray emitting particles by a factor of few hundred. This is in agreement with current models of PWNe spectra, which require $\\kappa\\simeq10^6$ if the radio band is included in the analysis \\citep{bucciantini_11}, a value which is two orders of magnitude larger than if only higher frequency bands are considered \\citep[$\\kappa\\simeq10^4$,][]{kennel_coroniti_84b}.  Although our work focuses primarily on the physics of PWNe, shocks propagating into striped flows may be present also in blazar jets and gamma-ray bursts \\citep[e.g.,][]{thompson_06}, where they may provide an appealing explanation, based on strong physical grounds, for any flat particle distribution inferred from observations. The broad-band emission of hotspots in radio galaxies, usually interpreted as the termination shocks of relativistic jets, indicates that the spectra of accelerated electrons need to be flat ($1<p<2$) below GeV energies \\citep{stawarz_07}. In the case of luminous blazar sources, very flat electron spectra below GeV energies are inferred directly from X-ray observations \\citep{celotti_08, sikora_09}. The steeper particle spectrum required at $\\gtrsim$GeV energies, with slope $p\\simeq2.5$, could instead result from Fermi acceleration across the shock, if the stripe-averaged field satisfies $|\\alpha|\\lesssim0.1$."
    },
    "1107/1107.2375_arXiv.txt": {
        "abstract": "We suggest that non-trivial correlations between the dark matter particle mass and collider based probes of missing transverse energy $H_{\\rm T}^{\\rm miss}$ may facilitate a two tiered approach to the initial discovery of supersymmetry and the subsequent reconstruction of the lightest supersymmetric particle (LSP) mass at the LHC.  These correlations are demonstrated via extensive Monte Carlo simulation of seventeen benchmark models, each sampled at five distinct LHC center-of-mass beam energies, spanning the parameter space of No-Scale $\\cal{F}$-$SU(5)$.  This construction is defined in turn by the union of the ${\\cal F}$-lipped $SU(5)$ Grand Unified Theory, two pairs of hypothetical TeV scale vector-like supersymmetric multiplets with origins in ${\\cal F}$-theory, and the dynamically established boundary conditions of No-Scale Supergravity. In addition, we consider a control sample comprised of a standard minimal Supergravity benchmark point. Led by a striking similarity between the $H_{\\rm T}^{\\rm miss}$ distribution and the familiar power spectrum of a black body radiator at various temperatures, we implement a broad empirical fit of our simulation against a Poisson distribution ans\\\"atz.  We advance the resulting fit as a theoretical blueprint for deducing the mass of the LSP, utilizing only the missing transverse energy in a statistical sampling of $\\ge$ 9 jet events. Cumulative uncertainties central to the method subsist at a satisfactory $12$-$15\\%$ level. The fact that supersymmetric particle spectrum of No-Scale $\\cal{F}$-$SU(5)$ has thrived the withering onslaught of early LHC data that is steadily decimating the Constrained Minimal Supersymmetric Standard Model and minimal Supergravity parameter spaces is a prime motivation for augmenting more conventional LSP search methodologies with the presently proposed alternative. ",
        "introduction": "\\subsection{Dark Matter and Missing Transverse Energy}   High-precision experiments have revealed that the matter content of our Universe is predominantly non-baryonic and dark.  The presence of large scale structure formation indicates that dark matter should be primarily cold, {\\it i.e.}~possessing a non-relativistic thermal speed at the time of its own decoupling.  Measurements of the anisotropies in the Cosmic Microwave Background (CMB) radiation, most recently and conclusively by the WMAP~\\cite{Spergel:2003cb,Spergel:2006hy,Komatsu:2010fb} satellite experiment, allocate the proportional energy densities of baryonic matter, cold dark matter (CDM), and dark energy (a cosmological constant) at about $4\\%$, $23\\%$, and $73\\%$ respectively. A stable Weakly Interacting Massive Particle (WIMP) with a canonical weak-scale mass can naturally have a thermal CDM relic density at the correct magnitude order~\\cite{Ellis:1983ew}. Although the Standard Model (SM) does not possess a CDM candidate, suitable particles are manifest in various SM extensions, most notably in constructions featuring supersymmetry (SUSY), which itself also efficiently addresses the quantum stability of the electroweak (EW) gauge hierarchy. Decay of the Lightest Supersymmetric Particle (LSP), typically associated with a neutralino eigenstate, may be blocked by a discrete $R$-parity symmetry, resulting in an exceedingly natural CDM candidate~\\cite{Ellis:1983ew,Goldberg:1983nd}. Alternative examples such as the little Higgs model with T-parity~\\cite{Cheng:2003ju, Cheng:2004yc} and extra dimensional models with KK parity~\\cite{Servant:2002aq} are not considered further here.  In general, if dark matter particles are produced in a collider, being $U(1)_{\\rm EM}$ neutral, $SU(3)_{\\rm C}$ singlets, they will escape the detector without direct observation.  The sole evidence of their production and rapid departure will be that of a deficit in the zero sum of momentum over all observed particles with a component trajectory which is perpendicular to the beamline.  This so-called ``missing transverse energy'' ($H_T^{\\rm miss}$) is measurable to fine precision at the LHC, which has been steadily accumulating data, thus far without any positively reported evidence for supersymmetry, from ${\\sqrt s}=7$ TeV proton-proton collisions since March 2010~\\cite{Khachatryan:2011tk, daCosta:2011hh, daCosta:2011qk}. In the current early operational stage of the LHC, our attention has unquestionably been piqued by the relentless severe reductions inflicted upon the viable minimal Supergravity (mSUGRA) and Constrained Minimal Supersymmetric Standard Model (CMSSM) parameter space.  It was concluded by one particular study that the miniscule $35~{\\rm pb}^{-1}$ of initially reported luminosity was alone sufficient to impeach more than $99\\%$ of the parameter space~\\cite{Strumia:2011}, although other analysts remained more circumspect~\\cite{Buchmueller:2011ki}.  The negative results have now held through the collection of more than a femtobarn of integrated luminosity by each of the ATLAS~\\cite{Aad:2011qa} and CMS~\\cite{PAS-SUS-11-003} collaborations, although we must remark that tantalizing low-statistics excesses do exist in both cases, which are neatly accounted for by simulation of the No-Scale $\\cal{F}$-$SU(5)$ collider-detector signal~\\cite{Li:2011av}.  Regardless, should this situation much longer persist into the deconstruction of the burgeoning wealth of collected raw data, reported to have already reached the impressive milestone of $5~{\\rm fb}^{-1}$ by the close of 2011, then the first great question of the LHC era shall become significantly heightened.  Namely, that is the question of whether there exist specifically detailed SUSY and/or superstring based post-SM constructions which simultaneously escape the existing experimental onslaught, yet also remain in principle discoverable to the present generation collider.  Granting, momentarily, an answer in the tacit affirmative to the prior question of discovery, one may speculate regarding the equally great question immediately subsequent, that being the issue of the detailed SUSY CDM classification. Differentiation between contending scenarios, both within the CMSSM and without, will require a probe at present and future collider and direct detection experiments, the latter now led by the XENON100 collaboration~\\cite{Aprile:2011hi}, of the intrinsic properties of the CDM, such as its spin and mass, and the structure of its interaction with SM particles. Assuming a SUSY LSP to represent the CDM, we may in principle make a collider based measurement of the LSP mass from the decay chains of heavier particles, for example linking the squark and gluino to the lightest neutralino in the supersymmetric SM~\\cite{Polesello:2004qy, Nojiri:2005ph}, though this method is highly model dependent, and there are myriad complications with its interpretation.  Firstly, since the momentum component along beamline can by its nature never be collected, it is impossible to reconstruct the true net missing energy total.  Secondly, $R$-parity implies that the LSP candidate will always ultimately appear pairwise.  Moreover, the two neutralinos may be kinematically decoupled from the original back-to-back pair production event, {\\it e.g.} of two heavy squarks $\\widetilde{q}\\,\\widetilde{\\overline{q}}$, by an intricate decay cascade, such that the final relative angular orientation is purely random, allowing for any possible partial cancellation. Finally, the dark matter component of the total missing energy $H_T^{\\rm miss}$ must be segregated from other potential sources, such as energetic neutrinos and energy mismeasurement, for example by particle escape through a calorimeter seam.  Although the prospect of a substantial loss of the latter type is unlikely, it competes with the ever more extraordinarily rare target events.  These foreboding circumstances compel some candid apprehension against the hope that any fundamental correlation might endure between the dark matter particle mass and the missing transverse energy measurement.  In the present work we set out to directly confront, by the aid of a massive Monte Carlo collider and detector simulation, the two great questions here posed.  The two tiered nature of our presentation, approaching first the prospect of SUSY CDM discovery, and secondly the subsequent classification of the SUSY CDM signal via explicit numerical correlation of the LSP mass with the peak of the missing transverse energy distribution, is in correspondence with our view of how the experimental reality might plausibly unfold in the coming months and years. Our study takes the context of a specific model, dubbed No-Scale $\\cal{F}$-$SU(5)$~\\cite{ Li:2010ws, Li:2010mi,Li:2010uu,Li:2011dw, Li:2011hr, Maxin:2011hy,Li:2011xu, Li:2011in,Li:2011rp,Li:2011fu,Li:2011ex,Li:2011av,Li:2011ab}, representing the merger of the ${\\cal F}$-lipped $SU(5)$ Grand Unified Theory (GUT)~\\cite{Barr:1981qv,Derendinger:1983aj,Antoniadis:1987dx}, two pairs of hypothetical TeV scale vector-like supersymmetric multiplets with origins in ${\\cal F}$-theory~\\cite{Jiang:2006hf,Jiang:2009zza,Jiang:2009za,Li:2010dp,Li:2010rz}, and the dynamically established boundary conditions of No-Scale Supergravity~\\cite{Cremmer:1983bf,Ellis:1983sf, Ellis:1983ei, Ellis:1984bm, Lahanas:1986uc}. The construction features an essential phenomenological consistency, profound predictive capacity, a singularly distinctive experimental signature, and imminent testability.  The core of the model space remains experimentally unblemished, in stark contrast to the incisive cuts already made into the heart of the more traditional mSUGRA/CMSSM formulations. This deft evasion of all existing LEP, Tevatron, and LHC sparticle mass limits stems from the characteristic sparticle mass hierarchy $m_{\\tilde{t}_1} < m_{\\tilde{g}} < m_{\\tilde{q}}$ of a light stop and gluino, both comfortably lighter than all other squarks, a feature which is stable across the full model space. If the tide advancing against the CMSSM should continue unabated, then island refuges such as No-Scale $\\cal{F}$-$SU(5)$ will demand increasing deference in the search for a physically feasible supersymmetric grand unified theory.   \\subsection{A No-Scale ${\\cal F}$-$SU(5)$ Primer}   The No-Scale $\\cal{F}$-$SU(5)$ construction inherits all of the most beneficial phenomenology~\\cite{Nanopoulos:2002qk} of flipped $SU(5)$~\\cite{Barr:1981qv,Derendinger:1983aj,Antoniadis:1987dx}, including fundamental GUT scale Higgs representations (not adjoints), natural doublet-triplet splitting, suppression of dimension-five proton decay and a two-step see-saw mechanism for neutrino masses, as well as all of the most beneficial theoretical motivation of No-Scale Supergravity~\\cite{Cremmer:1983bf,Ellis:1983sf, Ellis:1983ei, Ellis:1984bm, Lahanas:1986uc}, including a deep connection to the string theory infrared limit (via compactification of the weakly coupled heterotic theory~\\cite{Witten:1985xb} or M-theory on $S^1/Z_2$ at the leading order~\\cite{Li:1997sk}), the natural incorporation of general coordinate invariance (general relativity), a mechanism for SUSY breaking which preserves a vanishing cosmological constant at the tree level (facilitating the observed longevity and cosmological flatness of our Universe~\\cite{Cremmer:1983bf}), natural suppression of CP violation and flavor-changing neutral currents, dynamic stabilization of the compactified spacetime by minimization of the loop-corrected scalar potential and a dramatic reduction in parameterization freedom.  Written in full, the gauge group of flipped $SU(5)$ is $SU(5)\\times U(1)_{X}$, which can be embedded into $SO(10)$. The generator $U(1)_{Y'}$ is defined for fundamental five-plets as $-1/3$ for the triplet members, and $+1/2$ for the doublet. The hypercharge is given by $Q_{Y}=( Q_{X}-Q_{Y'})/5$.  There are three families of Standard Model (SM) fermions, whose quantum numbers under the $SU(5)\\times U(1)_{X}$ gauge group are \\begin{equation} F_i={\\mathbf{(10, 1)}} \\quad;\\quad {\\bar f}_i={\\mathbf{(\\bar 5, -3)}} \\quad;\\quad {\\bar l}_i={\\mathbf{(1, 5)}}, \\label{eq:smfermions} \\end{equation} where $i=1, 2, 3$.  There is a pair of ten-plet Higgs for breaking the GUT symmetry, and a pair of five-plet Higgs for electroweak symmetry breaking (EWSB). \\begin{eqnarray} & H={\\mathbf{(10, 1)}}\\quad;\\quad~{\\overline{H}}={\\mathbf{({\\overline{10}}, -1)}} & \\nonumber \\\\ & h={\\mathbf{(5, -2)}}\\quad;\\quad~{\\overline h}={\\mathbf{({\\bar {5}}, 2)}} & \\label{eq:Higgs} \\end{eqnarray} Since we do not observe mass degenerate superpartners for the known SM fields, SUSY must itself be broken around the TeV scale. In the minimal supergravities (mSUGRA), this occurs first in a hidden sector, and the secondary propagation by gravitational interactions into the observable sector is parameterized by universal SUSY-breaking ``soft terms'' which include the gaugino mass $M_{1/2}$, scalar mass $M_0$ and the trilinear coupling $A$.  The ratio of the low energy Higgs vacuum expectation values (VEVs) $\\tan \\beta$, and the sign of the SUSY-preserving Higgs bilinear mass term $\\mu$ are also undetermined, while the magnitude of the $\\mu$ term and its bilinear soft term $B_{\\mu}$ are determined by the $Z$-boson mass $M_Z$ and $\\tan \\beta$ after EWSB.  In the simplest No-Scale scenario, $M_0$=A=$B_{\\mu}$=0 at the unification boundary, while the complete collection of low energy SUSY breaking soft-terms evolve down with a single non-zero parameter $M_{1/2}$.  Consequently, the particle spectrum will be proportional to $M_{1/2}$ at leading order, rendering the bulk ``internal'' physical properties invariant under an overall rescaling. The rescaling symmetry can likewise be broken to a certain degree by the vector-like mass parameter $M_{\\rm V}$, although this effect is weak.  The matching condition between the low-energy value of $B_\\mu$ that is demanded by EWSB and the high-energy $B_\\mu = 0$ boundary is notoriously difficult to reconcile under the renormalization group equation (RGE) running.  The present solution relies on modifications to the $\\beta$-function coefficients that are generated by the inclusion of the extra vector-like multiplets, which may actively participate in radiative loops above their characteristic mass threshold $M_{\\rm V}$. Naturalness in view of the gauge hierarchy and $\\mu$ problems suggests that the mass $M_{\\rm V}$ should be of the TeV order. Avoiding a Landau pole for the strong coupling constant restricts the set of vector-like multiplets which may be given a mass in this range to only two constructions with flipped charge assignments, which have been explicitly realized in the $F$-theory model building context~\\cite{Jiang:2006hf,Jiang:2009zza, Jiang:2009za}.  We adopt the multiplets \\begin{eqnarray} & {XF} \\equiv {\\mathbf{(10,1)}} = (XQ,XD^c,XN^c) ~;~ {\\overline{XF}} \\equiv {\\mathbf{({\\overline{10}},-1)}} & \\nonumber \\\\ & {Xl} = {\\mathbf{(1, -5)}} \\quad;\\quad{\\overline{Xl}} = {\\mathbf{(1, 5)}}\\equiv XE^c \\, , & \\label{eq:flippons} \\end{eqnarray} where $XQ$, $XD^c$, $XE^c$ and $XN^c$ carry the same quantum numbers as the quark doublet, right-handed down-type quark, charged lepton and neutrino, respectively.  Alternatively, the pair of $SU(5)$ singlets may be discarded, but phenomenological consistency then requires the substantial application of unspecified GUT thresholds.  In either case, the (formerly negative) one-loop $\\beta$-function coefficient of the strong coupling $\\alpha_3$ becomes precisely zero, flattening the RGE running, and generating a wide gap between the large $\\alpha_{32} \\simeq \\alpha_3(M_{\\rm Z}) \\simeq 0.11$ and the much smaller $\\alpha_{\\rm X}$ at the scale $M_{32}$ of the intermediate flipped $SU(5)$ unification of the $SU(3) \\times SU(2)_{\\rm L}$ subgroup.  This facilitates a very significant secondary running phase up to the final $SU(5) \\times U(1)_{\\rm X}$ unification scale~\\cite{Li:2010dp}, which may be elevated by 2-3 orders of magnitude into adjacency with the Planck mass, where the $B_\\mu = 0$ boundary condition fits like hand to glove~\\cite{Ellis:2001kg,Ellis:2010jb,Li:2010ws}. This natural resolution of the ``little hierarchy'' problem corresponds also to true string-scale gauge coupling unification in the free fermionic string models~\\cite{Jiang:2006hf,Lopez:1992kg} or the decoupling scenario in F-theory models~\\cite{Jiang:2009zza,Jiang:2009za}, and also helps to address the monopole problem via hybrid inflation.  The modifications to the $\\beta$-function coefficients from introduction of the vector-like multiplets have a parallel effect on the RGEs of the gauginos.  In particular, the color-charged gaugino mass $M_{\\rm 3}$ likewise runs down flat from the high energy boundary, obeying the relation $M_3/M_{1/2} \\simeq \\alpha_3(M_{\\rm Z})/\\alpha_3(M_{32}) \\simeq \\mathcal{O}\\,(1)$, which precipitates a conspicuously light gluino mass assignment. The $SU(2)_{\\rm L}$ and hypercharge $U(1)_{\\rm Y}$ associated gaugino masses are by contrast driven downward from the $M_{1/2}$ boundary value by roughly the ratio of their corresponding gauge couplings $(\\alpha_2,\\alpha_{\\rm Y})$ to the strong coupling $\\alpha_{\\rm s}$. The large mass splitting expected from the heaviness of the top quark via its strong coupling to the Higgs (which is also key to generating an appreciable radiative Higgs mass shift $\\Delta~m_h^2$~\\cite{Li:2011ab}) is responsible for a rather light stop squark $\\widetilde{t}_1$. The distinctively predictive $m_{\\tilde{t}_1} < m_{\\tilde{g}} < m_{\\tilde{q}}$ mass hierarchy of a light stop and gluino, both much lighter than all other squarks, is stable across the full No-Scale $\\cal{F}$-$SU(5)$ model space, but is not precisely replicated in any phenomenologically favored constrained MSSM (CMSSM) constructions of which we are aware.  This spectrum generates a unique event topology starting from the pair production of heavy squarks $\\widetilde{q} \\widetilde{\\overline{q}}$, except for the light stop, in the initial hard scattering process, with each squark likely to yield a quark-gluino pair $\\widetilde{q} \\rightarrow q \\widetilde{g}$.  Each gluino may be expected to produce events with a high multiplicity of virtual stops, via the (possibly off-shell) $\\widetilde{g} \\rightarrow \\widetilde{t}$ transition, which in turn may terminate into hard scattering products such as $\\rightarrow W^{+}W^{-} b \\overline{b} \\widetilde{\\chi}_1^{0}$ and $W^{-} b \\overline{b} \\tau^{+} \\nu_{\\tau} \\widetilde{\\chi}_1^{0}$, where the $W$ bosons will produce mostly hadronic jets and some leptons. The model described may then consistently exhibit a net product of eight or more hard jets emergent from a single squark pair production event, passing through a single intermediate gluino pair, resulting after fragmentation in a spectacular signal of ultra-high multiplicity final state jet events. We remark also that the entirety of the viable $\\cal{F}$-$SU(5)$ parameter space naturally features a dominantly Bino LSP, at a purity greater than 99.7\\%, as is exceedingly suitable for direct detection, for example by XENON100~\\cite{Aprile:2011hi,Li:2011in}.  There exists no direct Bino to Wino mass mixing term. This distinctive and desirable model characteristic is guaranteed by the relative heaviness of the Higgs bilinear mass $\\mu$, which in the present construction generically traces the universal gaugino mass $M_{1/2}$ at the boundary scale $M_{\\cal F}$, and subsequently transmutes under the RGEs to a somewhat larger value at the electroweak scale.  A majority of the bare-minimally constrained~\\cite{Li:2011xu} parameter space of No-Scale $\\cal{F}$-$SU(5)$, as defined by consistency with the world average top-quark mass $m_{\\rm t}$, the No-Scale boundary conditions, radiative EWSB, the centrally observed WMAP7 CDM relic density~\\cite{Komatsu:2010fb}, and precision LEP constraints on the lightest CP-even Higgs boson $m_{h}$~\\cite{Barate:2003sz,Yao:2006px} and other light SUSY chargino and neutralino mass content, remains viable even after careful comparison against the first inverse femtobarn of LHC data~\\cite{Li:2011fu,Li:2011av}. Moreover, a highly favorable ``golden'' subspace~\\cite{Li:2010ws,Li:2010mi,Li:2011xg} exists which may simultaneously account for the key rare process limits on the muon anomalous magnetic moment $(g~-~2)_\\mu$ and the branching ratio of the flavor-changing neutral current decays $b \\to s\\gamma$ and $B_{s}^{0} \\to \\mu^+\\mu^-$.  The intersection of these experimental bounds is highly non-trivial, as the tight theoretical constraints, most notably the vanishing of $B_\\mu$ at the high scale boundary, render the residual parameterization deeply insufficient for arbitrary tuning of even isolated predictions, let alone the union of all predictions. In addition, a top-down consistency condition on the gaugino boundary mass $M_{1/2}$, and the parametrically coupled value of $\\tan\\beta$, is dynamically determined at a secondary local minimization $dV_{\\rm min}/dM_{1/2}=0$ of the minimum of the Higgs potential $V_{\\rm min}$; since $M_{1/2}$ is related to the modulus field of the internal string theoretic space, this also represents a dynamic stabilization of the compactification modulus.  The result is demonstrably consistent with the bottom-up phenomenological approach ~\\cite{Li:2010uu,Li:2011dw,Li:2011xu,Li:2011ex}, and we note in particular a rather distinctive conclusion that is enforced in both perspectives: the ratio $\\tan\\beta$ must have a value very close to 20. ",
        "conclusions": "The LHC era is rapidly evolving, with each major experimental collaboration having already accumulated a reported $5~{\\rm fb}^{-1}$ of integrated luminosity at the close of 2011.  Looking into the not so distant future, we must plan to focus not only on initial attainment of the much anticipated SUSY signal discovery, but also on mining the wealth of information present in the data for clues to the reconstruction of a detailed SUSY mass spectrum. We have concentrated in this presentation on precisely such a two tiered objective. Firstly, we have comprehensively documented the prospects for discovery of the full No-Scale $\\cal{F}$-$SU(5)$ model space at various LHC center-of-mass energies, based on the raw count of events with greater than $100$~GeV of missing transverse energy and at least nine jets. Secondly, upon observing the dual curiosities of i) a material correlation between the $H_{\\rm T}^{\\rm miss}$ histogram peak and a small multiple of the LSP mass, and ii) the strikingly familiar shape of the missing energy distribution, reminiscent of the power spectrum of a black body radiator at various temperatures, we have implemented a broad empirical fit of our extensive Monte Carlo $H_{\\rm T}^{\\rm miss}$ simulation against a Poisson distribution ans\\\"atz.  We advanced the resulting fit as a theoretical blueprint for deducing the mass of the LSP, utilizing only the missing transverse energy in a statistical distribution of $\\ge$ 9 jet events.  We judge the error directly associated with this technique to be within the satisfactory range of $12-15\\%$.  It will be interesting to see whether this supersymmetry mass extraction instrument might find some value by inclusion in the experimentalist toolbag, as we journey into and beyond the presumed discovery phases of the LHC operation in the next two years.  We have presented this analysis in the context of the No-Scale ${\\cal F}$-$SU(5)$ model, distinguished by its essential phenomenological consistency, profound predictive capacity, singularly distinctive experimental signature, and imminent testability.  The core of the model space around and above $M_{1/2} = 500$~GeV and $m_{\\rm LSP} = 100$~GeV remains experimentally unblemished, in stark contrast to the incisive cuts already made into the heart of the more traditional mSUGRA/CMSSM formulations.  This deft evasion of all existing LEP, Tevatron, and LHC sparticle mass limits stems from the characteristic sparticle mass hierarchy $m_{\\tilde{t}_1} < m_{\\tilde{g}} < m_{\\tilde{q}}$ of a light stop and gluino, both comfortably lighter than all other squarks, a feature which is stable across the full model space.  With some reports claiming that less than one percent of the formerly viable mSUGRA and CMSSM parameter space has withstood even the earliest accumulation of LHC data, the existence of a clearly defined, deeply theoretically motivated model which simultaneously exhibits consistency with all current experiments alongside testability at near term future experiments, is of some pressing interest.  The seventeen example spectra which we have studied span the model space of bare minimal constraints on No-Scale $\\cal{F}$-$SU(5)$, and establish a material correlation between the mass of the dark matter particle, taken to be the LSP neutralino, and the collider missing energy signal.  It appears that the stability of the SUSY mass hierarchy, with the overall scale driven by a single dimensionful parameter at the GUT scale, may be responsible for facilitating the remarkably robust and simple fitting relationship which has been observed.  Together, these facts speak efficiently to the suitability of No-Scale $\\cal{F}$-$SU(5)$ as a leading contender in the race to describe post Standard Model physics at the dawn of the LHC era."
    },
    "1107/1107.0831_arXiv.txt": {
        "abstract": "\\noindent As a formation route for objects such as giant planets and low-mass stars in protostellar discs (as well as stars in AGN discs), theories of self-gravitating disc fragmentation need to be able to predict the initial masses of fragments.  We describe a means by which the local Jeans mass inside the spiral structure of a self-gravitating disc can be estimated.  If such a self-gravitating disc satisfies the criteria for disc fragmentation, this estimate provides a lower limit for the initial mass of any fragments formed. We apply this approach to a series of self-gravitating protostellar disc models, to map out the typical masses of fragments produced by this formation mode.  We find a minimum fragment mass of around 3 Jupiter masses, which is insensitive to the stellar mass, and that - within the parameter space surveyed - fragments with masses between 10 and 20 Jupiter masses are the most common.  We also describe how the Jeans mass allows us to derive a more general criterion for disc fragmentation, which accounts for the processes of viscous heating, radiative cooling, accretion and the disc's thermal history.  We demonstrate how such a criterion can be determined, and show that in limiting cases it recovers several fragmentation criteria that have been posited in the past, including the minimum cooling time/maximum stress criterion. ",
        "introduction": "\\noindent During the process of low mass star formation, an excess of angular momentum in the parent molecular cloud will generally result in the formation of a circumstellar disc.  In the early phases of this disc's existence, it is expected to be relatively massive, cool, and extremely weakly ionised.  Therefore, disc self-gravity is expected to play a dominant role in protostellar and protoplanetary disc evolution \\citep{Lin1987,Laughlin1994}, in the same manner it is expected to play a role in the evolution of Active Galactic Nuclei (AGN), which also possess accretion discs \\citep{Shlosman1989,Goodman2003}.  A self-gravitating disc becomes gravitationally unstable if \\citep{Toomre_1964}:  \\begin{equation} Q = \\frac{c_s \\kappa}{\\pi G \\Sigma} \\sim 1, \\end{equation}  \\noindent where $c_s$ is the local sound speed, $\\Sigma$ is the surface density and $\\kappa$ is the epicyclic frequency (for Keplerian discs, this is equal to the angular frequency $\\Omega$).  This result applies primarily to axisymmetric perturbations, and is usually extended to non-axisymmetric perturbations by changing the critical $Q$ value to $Q< 1.5-1.7$, a result established empirically by numerical simulations (see \\citealt{Durisen_review} for a review).  The onset of gravitational instability (GI) produces non-axisymmetric, spiral structures.  It has been shown that self-gravitating discs will typically settle into what is known as a marginally stable state, where the stress heating produced by the spiral waves is balanced (at least approximately) by radiative cooling \\citep{Paczynski1978}.  In the marginally stable state, the spiral waves become quasi-steady structures that, while individually transient, as a whole remain for long timescales, mediating the transport of angular momentum outward (and the transport of mass inward).  The evolution of self-gravitating discs can therefore be neatly parametrised using the phenomenology of viscosity, much in the same way that the action of magneto-rotational instability (MRI) has been parametrised in less massive, more ionised discs \\citep{Blaes1994,Balbus1999}.  This pseudo-viscous parametrisation, and its encapsulation in the dimensionless $\\alpha$ parameter \\citep{Shakura_Sunyaev_73}, has been a useful tool in simplifying the computation of the structure of self-gravitating discs \\citep{Rafikov_05,Clarke_09,Rice_and_Armitage_09}, as well as computing the evolution of said discs \\citep{Lynden-Bell1974,Pringle1981a,Lin1990,Rice2010}.  While useful, this pseudo-viscous approximation is precisely that - an approximation, which is accurate only for a limited region of the available disc parameter space.  In particular, for the approximation to be valid, the angular momentum transport that occurs must be local.  Gravity is an inherently non-local force, and if the spiral structures in the disc can exert sufficient gravitational field stresses on fluid elements at large separations, then the pseudo-viscous approximation begins to fail, and the disc's angular momentum transport is governed by global properties \\citep{Balbus1999}.  The breakdown of the local approximation has been shown to occur when the disc-to-star mass ratio, $q$, is greater than 0.5, and/or the disc scale height satisfies $H/r > 0.1$ \\citep{Lodato_and_Rice_04,Forgan2011}.  Discs which satisfy these conditions typically possess strong $m=2$ spiral structures, which exert significant torques from large distances, and exhibit highly variable temperature structures, and consequently highly variable $Q$.  We will see that the issue of whether disc angular momentum transport is local or global will come to be of great importance in this paper.  If the transport is indeed local, the balance between the magnitude of disc stresses and the magnitude of radiative cooling has also been useful in describing the fragmentation of local self-gravitating discs.  \\citet{Gammie} demonstrated that a second criterion, alongside the $Q$-criterion above, must be satisfied if a self-gravitating disc is to fragment.  This criterion is related to the local cooling time, $t_{cool}$, and is usually expressed using the dimensionless cooling time parameter $\\beta_{c}=t_{cool} \\Omega$, i.e.  \\begin{equation} \\beta_{c} \\leq \\beta_{crit}, \\end{equation}  \\noindent where the exact value of $\\beta_{crit}$ is typically established by numerical simulations.  Using the 2D shearing sheet approximation, \\citet{Gammie} showed that when the ratio of specific heats in two dimensions $\\gamma_{2D}=2$, $\\beta_{crit} = 3$.  \\citet{Rice_et_al_05} extended this analysis to 3D smoothed particle hydrodynamics (SPH) simulations (where $\\beta_{c}$ was fixed at some value), and explored the dependence of $\\beta_{crit}$ on $\\gamma$.  They were able to confirm Gammie's assertion that, in the case where the disc has reached thermal equilibrium (i.e. the radiative cooling and stress heating due to gravitational instability are in balance), the fragmentation criterion can be recast in terms of the $\\alpha$ parameter:  \\begin{equation} \\alpha = \\frac{4}{9 \\gamma(\\gamma-1)\\beta_c}, \\end{equation}  \\noindent where the critical value of $\\alpha=\\alpha_{crit}\\approx 0.06$ for all values of $\\gamma$.  The fragmentation criterion can now be interpreted as a maximum GI stress that the disc can support in thermal equilibrium without fragmentation.  The inner regions of discs are typically too hot for $Q\\sim 1$, and are therefore not gravitationally unstable and incapable of fragmentation.  Also, if the primary contribution to $\\alpha$ is from GI stress, it will typically decrease with increasing proximity to the central star \\citep{Armitage_et_al_01,Rice_and_Armitage_09,Rice2010}.  These facts immediately suggest a minimum radius at which fragmentation can occur.  Many authors, using a variety of techniques, including semi-analytic methods, grid-based and particle based simulations, find that for self-gravitating protostellar discs the minimum is approximately $40$ au, using typical disc parameters \\citep{Rafikov_05, Matzner_Levin_05, Whit_Stam_06,Mejia_3,Stamatellos2008, intro_hybrid, Clarke_09}.  This also constrains the maximum available mass to form a fragment (as typically equilibrium surface density profiles decrease with increasing distance).  Despite such constraints, the consensus view is that the objects formed occupy the higher mass end of the planetary regime (and the low mass end of the stellar regime).  \\citet{Kratter2009} estimate the available mass within the most unstable wavelength to fragmentation (the disc scale height $H$) using one-dimensional disc models to obtain fragment masses, and then model the fragments' growth embedded in the disc to constrain expected planet masses.  Studies of the behaviour of $\\beta_c$ (or more strictly $t_{cool}$) in similar disc models \\citep{Rafikov_05, Nero2009} tend to agree with \\citet{Kratter2009} in that they calculate a minimum fragment mass of around $5 M_{\\rm Jup}$, where $M_{\\rm Jup}$ is the mass of Jupiter.  While the $Q$ and $\\beta_c$ criteria provide an elegant means by which to predict disc fragmentation, there are complicating factors that must also be considered.  Perhaps the most important is the fact that (for the most part) these numerically simulated discs are \\emph{isolated}.  We expect that self-gravitating discs - as they form at an early phase in protostellar evolution - will be surrounded by a substantial envelope from which the disc can accrete.  The potential for mass-loading as a result of accretion can alter the local optical depth as well as the local stresses felt by the disc.  \\citet{Kratter2010a} show in grid-based numerical simulations that systems accreting rapidly can sustain $\\alpha \\sim 1$ without fragmenting.  This is also confirmed by \\citet{Harsono2011} using SPH simulations, where they ascribe the extra Reynolds stresses to velocity shear at the disc surface.  Both show that the disc infall rate and disc-to-star accretion rate are typically proportional to each other, and that non-local angular momentum transport is significant.  These factors alone show the limited validity of isolated, constant $\\beta_c$ simulations in estimating disc fragmentation criteria.  Further to this, the imposition of a fixed $\\beta_c$ is somewhat unphysical.  As the disc cools through varying opacity regimes, the opacity is expected to change significantly (see e.g. \\citealt{intro_hybrid, Cossins2008}).  \\citet{Clarke2007} carry out SPH simulations where $\\beta_c$ is initially large, and gradually decreased towards the critical value.  They find that the eventual $\\beta_{crit}$ can be up to a factor of 2 lower.  This reduction of $\\beta_{crit}$ will increase $\\alpha_{crit}$ by the same factor - while the results of \\citet{Rice_et_al_05} provide a useful watermark, the exact value of $\\alpha_{crit}$ is not precisely determined.  Indeed, a related result of \\citet{Cossins2010} shows that $\\beta_{crit}$ also depends on the opacity regime the gas is in, producing isolated discs which require $\\alpha_{crit}=0.1$ for fragmentation.  Taking these factors into consideration, it becomes clear that the current use of $\\beta_c$ or $\\alpha$ is not sufficiently general as a fragmentation criterion.  While it is suitable for self-gravitating discs of constant $\\beta_c$ that do not accrete, such discs are not typically found in astrophysical situations.  Ideally, we would like to move towards a situation where the fragmentation criterion is sufficiently rich in its content that it can describe the full complexity of the fragmentation process in realistic self-gravitating discs.  This paper introduces two useful analytical formalisms: the first describes the local Jeans mass in a marginally stable, self-gravitating disc (which is described in section \\ref{sec:jeansmass}); this provides a means of predicting the masses of objects formed as a result of disc fragmentation.  In section \\ref{sec:betaj} we demonstrate that this expression for the Jeans mass can be used to develop a new timescale criterion, which folds in the influences of stress heating, radiative cooling, accretion and thermal history in general.  Having developed these concepts, we construct simple one-dimensional models to investigate the typical fragment masses produced by self-gravitating discs (described in section \\ref{sec:method}), the results and discussion of which can be seen in section \\ref{sec:results}.  We summarise the work in section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  \\noindent We have derived the local Jeans mass inside a spiral perturbation, and shown that it can be calculable using azimuthally averaged disc variables.  This allows the expected mass of objects formed from disc fragmentation to be calculated simply, both in theoretical models and in numerical simulations, across a variety of size scales.  We have taken this expression and derived a Jeans timescale $\\Gamma_J$, which describes the length of time required for the Jeans Mass to increase or decrease by its current value (normalised by the local angular velocity, in much the same way as the cooling time is often normalised).  We show that the resulting expression for $\\Gamma_J$ reduces to the standard cooling time criterion as was previously defined \\citep{Gammie,Rice_et_al_05}, and folds in related results on the disc's thermal history \\citep{Clarke2007} and envelope accretion \\citep{Kratter2010a}.  We confirm that subjecting the disc to extra shock heating promotes fragmentation, and that rapidly accretion encourages discs to fragment, provided that local angular momentum transport and cooling is efficient.  We have investigated the functional form of the Jeans Mass by applying it to simple one-dimensional steady-state self-gravitating disc models.  We find that the minimum fragment mass is around 3 Jupiter Masses, in agreement to previous work on the subject \\citep{Rafikov_05,Kratter2009,Nero2009}.  We also find that in a uniform survey of disc parameters, the most frequent fragment masses lie between 10 - 20 $M_{\\rm Jup}$.  These results are relatively insensitive to stellar mass, although lower mass stars require a higher disc-to-star mass ratio to produce fragments than their higher-mass counterparts.  Conversely, the mass of fragments found around lower mass stars are higher than that for higher mass stars (for the same outer radius and steady-state accretion rate).  We do however note that the derivation of the Jeans mass relies on assuming local angular momentum transport, which is often not the case for discs of high disc-to-star mass ratio.  The errors introduced by this assumption require further empirical data about non-local angular momentum transport to quantify accurately.  The formalisms introduced in this paper are amenable to a variety of studies in self-gravitating discs of all scales, allowing the intial fragment mass distribution to be identified given the known disc population, and to provide more general criteria for disc fragmentation itself."
    },
    "1107/1107.3206.txt": {
        "abstract": "We present a new approach for identifying the Tip of the Red Giant Branch (TRGB) which, as we show, works robustly even on sparsely populated targets. Moreover, the approach is highly adaptable to the available data for the stellar population under study, with prior information readily incorporable into the algorithm. The uncertainty in the derived distances is also made tangible and easily calculable from posterior probability distributions. We provide an outline of the development of the algorithm and present the results of tests designed to characterize its capabilities and limitations. We then apply the new algorithm to three M31 satellites: Andromeda I, Andromeda II and the fainter Andromeda XXIII, using data from the Pan-Andromeda Archaeological Survey (PAndAS), and derive their distances as $731^{(+ 5) + 18}_{(- 4) - 17}$ kpc, $634^{(+ 2) + 15}_{(- 2) - 14}$ kpc and $733^{(+ 13)+ 23}_{(- 11) - 22}$ kpc respectively, where the errors appearing in parentheses are the components intrinsic to the method, while the larger values give the errors after accounting for additional sources of error. These results agree well with the best distance determinations in the literature and provide the smallest uncertainties to date. This paper is an introduction to the workings and capabilities of our new approach in its basic form, while a follow-up paper shall make full use of the method's ability to incorporate priors and use the resulting algorithm to systematically obtain distances to all of M31's satellites identifiable in the PAndAS survey area. ",
        "introduction": "The Tip of the Red Giant Branch is a very useful standard candle for gauging distances to extended, metal-poor structures. The tip corresponds to the very brightest members of the First Ascent Red Giant Branch (RGB), at which point stars are on the brink of fusing helium into carbon in their cores and hence contracting and dimming to become Horizontal Branch stars. The result is a truncation to the Red Giant Branch when the Color-Magnitude Diagram (CMD) for an old stellar population is generated, beyond which lie only the comparatively rare Asymptotic Giant Branch Stars and sources external to the system of interest. The (highly variable) contamination from such objects provides the principal obstacle to simply `reading off' the tip position from the RGB's luminosity function and the truncation of the AGB can even masquerade as the TRGB in certain instances. The I-band is the traditionally favored region of the spectrum for TRGB measurements, minimizing the interstellar reddening that plagues shorter wavelengths, while keeping dependance on metallicity lower than it would be at longer IR wavelengths. It should also be remembered that stars approaching the TRGB generally exhibit peak emission in this regime. \\citet{Iben83} determined that low mass ($< 1.6 \\, {\\rm M}_{\\odot}$ for Pop. I, $< 1 \\, {\\rm M}_{\\odot}$ for Pop. II), metal-poor ([Fe/H] $<$ -0.7 dex) stars older than 2 Gyr produce a TRGB magnitude that varies by only 0.1 magnitudes. More recently, \\citet{Bellazzini01} determined the tip magnitude to lie at an I-band magnitude of $M_{TRGB} = -4.04 \\pm 0.12$. This low variation can be attributed to the fact that all such stars have a degenerate core at the onset of helium ignition and so their cores have similar properties regardless of the global properties of the stars. The result is a standard candle that is widely applicable to the old, metal-poor structures that occupy the halos of major galaxies. Distances derived from the TRGB, unlike those from the Cepheid variable or RR Lyrae star for example, can be determined from a single epoch of observation, making it very useful for wide-area survey data. Furthermore, \\citet{Salaris97} confirmed agreement between Cepheid and RR Lyrae distances and TRGB distances to within $\\sim$5 \\%.  Up until \\citet{Lee93} published their edge-finding algorithm, the tip had always been found by eye, but clearly if the wide-reaching applications of the TRGB standard candle were to be realized, a more consistent, repeatable approach was in order. The aforementioned paper shows that, if a binned luminosity function (LF) for the desired field is convolved with a zero sum Sobel kernel [-2, 0, +2], a maximum is produced at the magnitude bin corresponding to the greatest discontinuity in star counts, which they attribute to the tip. Using this method, they were able to obtain accuracies of better than 0.2 magnitudes. \\citet{Sakai96} set out to improve on this approach by replacing the binned LF and kernel with their smoothed equivalents. To do this, they equate each star with a Gaussian probability distribution whose FWHM is determined by the photometric error at the magnitude actually recorded for the star. Then rather than each star having to fall in a single bin, it contributes to all bins, but most strongly to the bin at the magnitude recorded for it and less so the further a bin is from that recorded magnitude, weighted by the photometric error. This is illustrated in equation \\ref{e_Sakai}:  \\begin{equation} \\Phi(m) = \\sum^N_{i=1}\\frac{1}{\\sqrt{2\\pi\\sigma^2_i}}\\exp{\\left[-\\frac{(m_i - m)^2}{2\\sigma^2_i}\\right]} \\label{e_Sakai} \\end{equation} where $m$ is the magnitude of the bin in question and $m_i$ and $\\sigma^2_i$ are the central magnitude and variance respectively of the gaussian probability distribution for the \\emph{i}th star. This method halved the error associated with the non-smoothed version of the algorithm and an identical smoothing is hence justly incorporated into the model LF for our Bayesian approach.  In a more recent variation on the Edge Detection methods, \\citet{Madore09} once again applied a Sobel kernel, but fit to a luminosity function built from composite stellar magnitudes $T \\equiv I - \\beta [(V-I)_0 - 1.50]$ where $\\beta$ is the slope of the TRGB as a function of color. This they argue results in a sharper output response from the filter, and allows all stars, regardless of color, to contribute equally to the derived tip position. \\citet{Rizzi07} derived a value of $0.22 \\pm 0.02$ for $\\beta$ after a study of 5 nearby galaxies, and show that it is quite consistent from one galaxy to another.  \\citet{Mendez02} made a departure from the simple `edge-finding' algorithms above by adapting a maximum likelihood model fitting procedure into their technique. They point out that the luminosity function faint-ward of the tip is well modeled as a power law:  \\begin{equation} L (m \\ge m_{TRGB}) = 10^{a(m - m_{TRGB})} \\label{e_LF} \\end{equation} where $m \\geq m_{TRGB}$ and $a$ is fixed at 0.3.  They then ascribe the location of the tip to the magnitude at which this power law truncates - i.e. $m = m_{TRGB}$.  Bright-ward of the tip they assume a functional form:  \\begin{equation} L (m < m_{TRGB}) = 10^{b(m - m_{TRGB})-c} \\end{equation} where $b$ is the slope of the power law bright-ward of the tip and $c$ is the magnitude of the step at the RGB tip.  Such a model, though simplistic, is robust against the strong Poisson noise that is inevitable in more sparsely populated LFs, making it a significant improvement over the previous, purely `edge-finding' methods.  \\citet{Makarov06} follow in a similar vein, demonstrating the proven advantages of a Maximum Likelihood approach over simple Edge Detection techniques, despite a model dependance. Unlike \\citet{Mendez02} however, they allow $a$ as a free parameter, arguing its notable variance from $0.3$ and importantly, they smooth their model LF using a photometric error function deduced from artificial star experiments. One shortcoming of both of these methods however, is that the most likely parameter values alone are obtained, without their respective distributions or representation of their dependance on the other parameters. Also, with regard to the background contamination, the RGB LF in fact sits on top of non-system stars in the field and so rather than model the background exclusively bright-ward of the TRGB, the truncated power law of Eq. \\ref{e_LF} can be added onto some predefined function of the contamination.  Arguably the most successful method developed so far has been that devised by \\citet{McConn04}. It has been used to ascertain accurate distances to 17 members of the Local Group \\citep{McConn05}. It combines aspects of both `edge-finding' and model fitting to zero in more accurately on the tip. They argue that as the precise shape of the luminosity function at the location of the tip is not known, a simple Sobel Kernel approach that assumes a sharp edge to the RGB, does not necessarily produce a maximum at the right location. They instead use a least-squares model fitting technique that fits to the LF in small windows searching for the portion best modeled by a simple slope function. This, they reason, marks the location of the steepest decline in star counts which is attributable to the tip location. This method is capable of finding the tip location accurate to better than 0.05 magnitudes, although is still susceptible to be thrown off by noise spikes in a poorly populated LF.  Despite the merits of previous methods such as these, none of them work particularly well when confronted with the high levels of Poisson noise that abound in the more poorly populated structures of galaxy halos. Furthermore, in such conditions as these where the offset between detected and true tip position will likely be at its greatest, it is of great use to have a full picture of likelihood space, as opposed to merely the determined, most probable value. This has lead us to develop a new, Bayesian approach to locating the TRGB, specifically, one that incorporates a Markov Chain Monte Carlo (MCMC) algorithm. As shall become apparent in the next section, such a method is very robust against noise spikes in the luminosity function and allows all prior knowledge about the system to be incorporated into the tip-finding process - something lacking in the previous approaches. Further to this, the MCMC provides for a remarkably simple, yet highly accurate error analysis. It also makes it possible to marginalize over parameters to provide posterior probability distributions (PPDs) of each parameter, or to obtain plots of the dependance of each parameter on every other. In \\S \\ref{s_MCMC_Method}, a detailed explanation of our approach and its limitations is given. \\S \\ref{ss_ANDI} introduces the method by applying the algorithm to one of M31's brightest dwarf spheroidals, Andromeda I. \\S \\ref{ss_errors} discusses the nature of systematic errors that apply to the method. \\S \\ref{ss_Initial_Tests} investigates the accuracy that the basic method (before addition of priors) is capable of given the number of stars populating the LF for the field and the strength of the non-RGB background while \\S \\ref{ss_Composite_LFs} deals with its performance when faced with a composite luminosity function. \\S \\ref{s_More_Distances} then applies our new approach to two additional M31 dwarf satellite galaxies and \\S \\ref{s_Conclusions} summarizes the advantages of the method and outlines the expected applicability of the method in the immediate future. ",
        "conclusions": "\\label{s_Conclusions}  The versatility and robustness of our new method can be appreciated from \\S \\ref{s_MCMC_Method} and its high level of accuracy is evident from the measurement errors which are consistently smaller than those in the literature to date. In addition, it is our hope that with the correct priors imposed, this new approach carries with it the ability to gauge distances to even the most poorly populated substructures, bringing a whole new range of objects with in reach of the TRGB standard candle. In the case of the M31 halo alone, it will be possible to obtain distances to all of the new satellites discovered by the PAndAS survey - a feat previously impractical using the TRGB. Furthermore, PAndAS has revealed a complicated network of tidal streams that contain valuable information as to the distribution of dark matter within the M31 halo. With our new method, it will be possible to systematically obtain distances at multiple points along these streams, thus providing vital information for constraining their orbits.  The great advantage of our new Bayesian method over a pure maximum likelihood method is the ease with which prior information may be built into the algorithm, making it more sensitive to the tip. Here in lies the great power of the Bayesian approach, whereby the addition of a few carefully chosen priors can reduce the measurement errors ten-fold. The result is an algorithm that is not only very accurate but highly adaptable and readily applicable to a wide range of structures within the distance (and metallicity) limitations of the TRGB standard candle. With instruments such as the 6.5 m Infra-red James Webb Space Telescope and the 42 m European Extremely Large Telescope expected to be operational within the decade, these distance limitations will soon be greatly reduced. This will bring an enormous volume of space within reach of the TRGB method, including the region of the Virgo Cluster. A tool with which it is possible to apply the TRGB standard candle to small, sparsely populated structures and small subsections of large structures alike is hence, needless to say, invaluable."
    },
    "1107/1107.5625_arXiv.txt": {
        "abstract": "{ To perform imaging observations of optically red objects such as high redshift quasars and brown dwarfs, the Center for the Exploration of the Origin of the Universe (CEOU) recently developed an optical CCD camera, Camera for QUasars in EArly uNiverse (CQUEAN), which is sensitive at 0.7-1.1 $\\mu$m. To enable observations with long exposures, we developed an auto-guiding system for CQUEAN. This system consists of an off-axis mirror, a baffle, a CCD camera, a motor and a differential decelerator. To increase the number of available guiding stars, we designed a rotating mechanism for the off-axis guiding camera. The guiding field can be scanned along the 10 arcmin ring offset from the optical axis of the telescope. Combined with the auto-guiding software of the McDonald Observatory, we confirmed that a stable image can be obtained with an exposure time as long as 1200 seconds.} ",
        "introduction": "Most telescopes require active guiding to stay on target for long periods. Atmospheric refraction, mechanical imperfections, flexure of the mount, and other errors make long unguided exposures difficult  \\citep{lee07}. Historically, astrophotographers guided by watching a star through an eyepiece while manually adjusting the position of the telescope or camera. Today guide cameras with charge-coupled device (CCD) sensors and computer control have revolutionized astronomy, automatically guiding  for long durations and with better accuracy than is humanly possible \\citep{bir06}.  In the auto-guiding system, an image of the guide star is taken repeatedly, and a computer calculates the centroids and commands the telescope control system (TCS) to compensate for drift.  On some instruments the guide star is the actual science target - for example the image on a spectrograph slit - but our guider captures the image from outside the science field with a separate movable guide camera. Another design uses a separate telescope for the guider.  In Korea, auto-guiding systems have been developed to improve tracking accuracy \\citep{jeo99, yoo06}, and \\citet{mun06} developed a telescope control program which has auto-guiding capability for the telescope in Kyung Hee University. Efforts to improve guiding accuracy \\citep{ise08} and pointing accuracy \\citep{kan06} have been made at other observatory. % Commercial auto-guiding programs are even popular among amateur astronomers.   \\begin{figure}[t] \\plotfiddle{Fig1.eps}{8cm}{270}{34}{44}{-140}{260} \\caption{Unguided drift of the 2.1m telescope along Right Ascension (top) and Declination (bottom) direction. It shows the 2-minute periodic error in RA caused by a slight misalignment of the final worm gear. The longer period motion is probably caused by slight imperfections in the gear teeth and flexure in the mount. High frequency motion in Dec may be caused by wind shake and oscillations of the mount that disturb the image centroid, and presumably the motion in RA contains the same kind of noise. The source of longer period motion in Dec is unknown.}% \\label{drift} \\end{figure}  CQUEAN is an optical camera system for the survey of high redshift ($z>5$) quasars in the early universe, which is developed by the Center for the Exploration of the Origin of the Universe (CEOU). It is composed of a science CCD camera, a focal reducer, seven broad-band filters (g', r', I', z', Y, Iz and Is) and an auto-guiding system. The science camera has a deep-depletion CCD whose sensitivity at 1um is better than conventional designs. It is attached to the 2.1m Otto Struve Telescope at McDonald Observatory, Texas, USA. The overall description of CQUEAN will be presented in \\citet{par11}.   In the high and dry desert of far West Texas, the McDonald Observatory is an excellent location for astronomical observations. The 2.1m Otto Struve Telescope is a classical cassegrain telescope built in 1939 by Warner \\& Swasey, and was the second largest telescope in the world. The cassegrain focal ratio is f/13.65. Like most telescopes of its era, it rotates on a equatorial mount.  An auto-guider was developed by \\citet{abb90}, but it was not successful. At present, one of the main instruments on the 2.1m telescope, the Sandiford Echelle Spectrograph, is now fully auto-guided, but the hardware and software is quite specific to the spectrograph. Thus we developed an auto-guiding system for CQUEAN in order to observe faint targets with long exposures.  In this paper, we describe the development of the CQUEAN auto-guiding system and analyze its guiding performance. In section 2, we show the tracking status of the 2.1m Otto Struve Telescope without auto-guiding, the sensitivity of the guide camera, and the expected number of guide stars for the estimated sensitivity. In section 3, the system design and the components of the auto-guiding system are described. In section 4, we present the analysis of the first observations, and make some concluding remarks in section 5. ",
        "conclusions": "We successfully developed, tested, and calibrated an auto-guiding system for CQUEAN. It consists of a camera, an off-axis mirror, and an arm structure to hold all components. To increase the expected number of guide stars that are brighter than the limiting magnitude, we built a mechanism that rotates the guide camera around the optical axis of the telescope. This is more efficient than adopting a larger format CCD chip fixed to a single position, because the 2.1 m Otto Struve Telescope is a classical cassegrain type which has strong coma aberration. The rotating mechanism minimizes the offset from the optical axis, and thus results in better PSFs in the field. From the astrometry on the guide CCD camera images, we estimated the offset amount caused by mechanical flexure of the arm and decelerator backlash. We found % that these factors do not significantly affect guiding performance. The comparison between two 600 second exposures made on the science camera with and without auto-guiding proves the auto-guider works well and is necessary. Exposures up to 1200 sec long were made with auto-guiding and the resultant images display round stellar profiles.%"
    },
    "1107/1107.2233_arXiv.txt": {
        "abstract": "Many hydrogen deficient stars are characterised by surface abundance patterns that are hard to reconcile with conventional stellar evolution. Instead, it has been suggested that they may represent the result of a merger episode between a helium and a carbon-oxygen white dwarf.  In this Letter, we present a nucleosynthesis study of the merger of a $0.4\\, M_{\\sun}$ helium white dwarf with a $0.8\\, M_{\\sun}$ carbon-oxygen white dwarf, by coupling the thermodynamic history of Smoothed Particle Hydrodynamics particles with a post-processing code.  The resulting chemical abundance pattern, particularly for oxygen and fluorine, is in qualitative agreement with the observed abundances in R Coronae Borealis stars. ",
        "introduction": "Hydrogen deficient stars with high carbon abundances (with the exceptions of the PG 1159 and CSPN WR stars) can be divided into three main categories according to their effective temperatures: Hydrogen Deficient Carbon (HdC) stars, R~Coronae~Borealis (RCB), and Extreme Helium (EHe) stars.  They are often considered to be at different evolutionary stages following a common origin, as evidenced by their similar atmospheric abundances \\citep{PAN04}.  The origin of these stars is of particular interest because of their unique chemical composition.  As well as being enriched in carbon and oxygen relative to solar abundances, they also exhibit enrichment in {nitrogen, neon, fluorine, and some s-processed material}. These abundance patterns are hard to reconcile with their expected initial composition, requiring alternative scenarios for their origin. Two leading theories are the Final Flash (FF) and the Double-Degenerate (DD) scenarios.  The FF scenario attributes the formation to a late He-shell flash in a post-Asymptotic Giant Branch (post-AGB) star as it evolves towards {the white dwarf state}. The flash would move the star back onto the AGB part of the Hertzsprung Russell diagram, while mixing s-processed material and He-burning ashes to the surface \\citep{FUJ77,REN79,WER06}.  In the DD scenario, the merger of two white dwarfs (He+CO), with a total mass below the Chandrasekhar limit, through disruption and accretion of the lighter star, gives rise to a single object whose composition is an admixture of the two white dwarfs \\citep{WEB84,ASP00}.  This dense object is expected to be surrounded by a hot corona and an accretion disk \\citep{LOR09}. Surface abundance analyses of hydrogen-deficient stars seem to favour the latter scenario\\ \\citep{SAI02}, although the possibility of some objects being formed through the FF scenario cannot be excluded \\citep{PAN11}.  When considering the formation of hydrogen deficient objects through the DD scenario, the following question can be raised: does nuclear processing occur during the merging event?  If no nuclear burning occurs --- the ``cold'' merger hypothesis --- the observed abundances might be reproduced through the {partial} mixing of material previously processed during the AGB phase of the CO white dwarf progenitor. However, recent calculations based on this scenario failed to qualitatively match the overall abundance pattern observed \\citep{JEF11}.  Another possibility involves a ``hot'' merger event, in which significant nuclear processing occurs.  Indeed, \\cite{CLA07} suggested that this scenario may account for the high abundances of $^{18}$O observed in some stars.  Unfortunately, no detailed nucleosynthesis study during white dwarf mergers has been published to date.  In this Letter, we present a detailed study of the nucleosynthesis accompanying white dwarf mergers. Calculations are based on post-processing analyses using temperature and density versus time profiles obtained from Smoothed Particle Hydrodynamics simulations of the merger event. The structure of the manuscript is as follows: in section \\ref{sec:observations} we summarise the observed abundances in hydrogen deficient objects. The details of previous nucleosynthesis studies and of the calculations presented in this work are given in sections \\ref{sec:previous-studies} and \\ref{sec:SPH}, respectively. Results are discussed in section \\ref{sec:nucleosynthesis}, and the main conclusions are reached in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Mergers of helium and carbon-oxygen white dwarfs are currently favoured as the events responsible for the production of hydrogen deficient HdC, EHe, and RCB stars. The main question facing us is whether these objects are the result of a cold merger, in which no nucleosynthesis occurs, or a hot merger, in which at least some of the material is processed through nucleosynthesis.  The results presented in this work show that nuclear processing during the merger event can account for some of the observed abundances in these stars, particularly for fluorine, which cannot be easily synthesised in the cold merger scenario.  There are, however, still open questions regarding the expected nucleosynthesis: neon is difficult to produce at the observed levels, while the absolute abundance of carbon is too high in our models. Such discrepancies may be alleviated by the inclusion of a more detailed initial chemical composition with many more species, resulting from state-of-the-art models of the preceding AGB phase.  Inclusion of such detailed abundances is currently under way."
    },
    "1107/1107.1154_arXiv.txt": {
        "abstract": "In the presence of large-scale vortical motions and/or magnetic-field strains, the turbulent cross helicity (velocity--magnetic-field correlation in fluctuations) may contribute to the turbulent electromotive force and the Reynolds stress. These effects of cross helicity are considered to balance the primary effects of turbulence such as the turbulent magnetic diffusivity in magnetic-field evolution and the eddy viscosity in the momentum transport. The cross-helicity effects may suppress the enhanced transports due to turbulence. Physical interpretation of the effects is presented with special emphasis on the difference between the cross-helicity effect and the usual $\\alpha$ or helicity effect in the dynamo action. The relative importance of the cross-helicity effect in dynamo action is validated with the aid of a direct numerical simulation (DNS) of the Kolmogorov flow with an imposed magnetic field. Several mechanisms that provide turbulence with the cross helicity are also discussed. ",
        "introduction": "Introduction} Cross helicity (velocity--magnetic-field correlation) is a quantity of fundamental importance at high magnetic Reynolds number magnetohydrodynamic (MHD) flow. The total cross helicity, as well as the MHD energy, is an inviscid invariant of the MHD system of equations, but a pseudoscalar which represents breakage of mirror-symmetry. Cross helicity is expected to play a certain role in the turbulent dynamo \\citep[see] [and works cited therein]{Yoshizawa_etal_2004}.  In the presence of a large-scale vortical motion, the turbulent cross helicity $\\langle {{\\boldsymbol{u}}' \\cdot {\\boldsymbol{b}}'} \\rangle$ contributes to the electromotive force aligned with the large-scale vorticity ${\\boldsymbol{\\it\\Omega}} (= \\nabla \\times {\\boldsymbol{U}})$ (${\\boldsymbol{u}}'$: velocity fluctuation, ${\\boldsymbol{b}}'$: magnetic-field fluctuation, ${\\boldsymbol{U}}$: mean velocity, $\\langle{\\cdots}\\rangle$: ensemble average). A fluid element in the large-scale vorticity is subject to the Coriolis-like force ${\\boldsymbol{u}}' \\times {\\boldsymbol{\\it\\Omega}}$. Provided that there is a positive (or negative) cross correlation between the velocity and magnetic-field fluctuations, a contribution to the electromotive force parallel (or antiparallel) to the mean vorticity, $\\langle {\\delta{\\boldsymbol{u}}' \\times {\\boldsymbol{b}}'} \\rangle \\propto \\langle {{\\boldsymbol{u}}' \\cdot {\\boldsymbol{b}}'} \\rangle {\\boldsymbol{\\it\\Omega}}$, arises (figure~\\ref{fig:cross-helicity_effect}). This is in marked contrast with a positive (or negative) turbulent helicity $\\langle {{\\boldsymbol{u}}' \\cdot {\\boldsymbol{\\it\\omega}}'} \\rangle$, correlation between the velocity and vorticity fluctuations may contribute to the electromotive force antiparallel (or parallel) to the mean magnetic field [${\\boldsymbol{\\it\\omega}}' (= \\nabla \\times {\\boldsymbol{u}}')$: vorticity fluctuation].  If we take the cross-helicity effects into account, the turbulent electromotive force ${\\boldsymbol{E}}_{\\rm{M}}$ is expressed as \\citep{Yoshizawa_1990} \\begin{equation} {\\boldsymbol{E}}_{\\rm{M}} \\equiv \\langle {{\\boldsymbol{u}}' \\times {\\boldsymbol{b}}'} \\rangle = \\alpha {\\boldsymbol{B}} - \\beta \\nabla \\times {\\boldsymbol{B}} + \\gamma {\\boldsymbol{\\it\\Omega}}, \\label{eq:turb_emf}% \\end{equation} where $\\alpha$, $\\beta$, and $\\gamma$ are the transport coefficients expressed in terms of the residual helicity $\\langle { - {\\boldsymbol{u}}' \\cdot {\\boldsymbol{\\it\\omega}}' + {\\boldsymbol{b}}'\\cdot {\\boldsymbol{j}}' } \\rangle (\\equiv H)$, the turbulent MHD energy $\\langle {{\\boldsymbol{u}}'{}^2 + {\\boldsymbol{b}}'{}^2} \\rangle/2 (\\equiv K)$, and the cross helicity $\\langle {{\\boldsymbol{u}}' \\cdot {\\boldsymbol{b}}' } \\rangle (\\equiv W)$, respectively. \\begin{figure}[htb] \\begin{center} \\includegraphics[width = 0.5\\textwidth] {yokoi_balarac_etc13_fig_01.eps} \\caption{\\label{fig:cross-helicity_effect}Physical interpretation of the cross-helicity effect. Redrawn from \\citet{Yokoi_1999}.} \\end{center} \\end{figure}  If we substitute (\\ref{eq:turb_emf}) into the mean magnetic induction equation, we obtain \\begin{equation} \\frac{\\partial {\\boldsymbol{B}}}{\\partial t} = \\nabla \\times \\left( { {\\boldsymbol{U}} \\times {\\boldsymbol{B}} } \\right) + \\nabla \\times \\left( { \\alpha {\\boldsymbol{B}} + \\gamma {\\boldsymbol{\\it\\Omega}} } \\right) - \\nabla \\times \\left[ { (\\eta + \\beta) \\nabla \\times {\\boldsymbol{B}} } \\right] \\label{eq:mean_mag_ind_eq}% \\end{equation} ($\\eta$: magnetic diffusivity). First, in the presence of turbulence, the effective magnetic diffusivity or resistivity is enhanced by turbulence ($\\eta \\rightarrow \\eta + \\beta$) with spatiotemporal variation of $\\beta$. So, the coefficient $\\beta$ is called the turbulent magnetic diffusivity or anomalous resistivity. Secondly, the $\\alpha$- and $\\gamma$-related terms express possible magnetic-field generation mechanisms due to pseudoscalars. If turbulence has some asymmetric properties represented by these pseudoscalars, the enhanced transport due to $\\beta$ can be suppressed by these pseudoscalar effects. This suppression of turbulent transport is based on the dynamic balance between two turbulent effects: enhancement and suppression of the transport. This state is entirely different from the one subject only to the molecular transport. Turbulent itself is very strong, but due to the additional asymmetry represented by pseudoscalars, effective turbulent transport is dynamically suppressed. Then important point to see is how and how much such pseudoscalars can be present in turbulence.  In the long history of the dynamo study, special attention has been paid to the $\\alpha$ or helicity effect [the first term in (\\ref{eq:turb_emf})]. In contrast, the $\\gamma$ or cross-helicity effect [the third term in (\\ref{eq:turb_emf})] has been almost missing except for some limited number of works \\citep[e.g.,][and works cited therein]{Yoshizawa_etal_2004}. One of the reasons of this missing may be attributed to the current treatment of turbulence; in most works, the large-scale velocity is neglected in the consideration of turbulence because of the Galilean invariance of the momentum equation. However, such treatment inevitably leads to the neglect of the large-scale shear effects. On the contrary, a large-scale rotation is ubiquitous and sometimes essential ingredient in astro/geophysical phenomena. In this sense, we should take effects of large-scale inhomogeneous flows into consideration of the realistic turbulent dynamo studies.  The number of works that contain numerical simulations examining or validating the cross-helicity effect is very limited \\citep{Hamba_1992,Nishino_Yokoi_1998,Yokoi_Hamba_2007,Yokoi_etal_2008,Sur_Brandenburg_2009}. In the present work, we examine the effect of turbulent cross helicity, and see the validity of the notion of cross-helicity dynamo. This must be a very interesting contribution to the astro/geophysical dynamo studies from the shear turbulence study.  Organization of this paper is as follows. After an outlined presentation of the turbulent cross-helicity generation mechanisms with the aid of the evolution equation of the turbulent cross helicity in section~\\ref{sec:level2}, a brief description of the MHD Kolmogorov flow is given in section~\\ref{sec:level3}. Some of the basic numerical results are presented in section~\\ref{sec:level4}, which include the analysis of the turbulent electromotive force (section~\\ref{sec:level4-1}), the comparison of the induced fields with some approximate dynamo solutions (section~\\ref{sec:level4-2}), the comparison of the spatial distributions of the turbulent cross helicity and its production rate (section~\\ref{sec:level4-3}), and the turbulent cross helicity scaled by the turbulent MHD energy (section~\\ref{sec:level4-4}). Concluding remarks with directions in the future work are presented in section~\\ref{sec:level5}. ",
        "conclusions": ""
    },
    "1107/1107.4547_arXiv.txt": {
        "abstract": "Here we present 1584 new southern proper motion systems with $\\mu$ $\\ge$ 0$\\farcs$18 yr$^{-1}$ and 16.5 $>$ $R_{59F}$ $\\ge$ 18.0. This search complements the six previous SuperCOSMOS-RECONS (SCR) proper motion searches of the southern sky for stars within the same proper motion range, but with $R_{59F}$ $\\le$ 16.5. As in previous papers, we present distance estimates for these systems and find that three systems are estimated to be within 25 pc, including one, SCR 1546-5534, possibly within the RECONS 10 pc horizon at 6.7 pc, making it the second nearest discovery of the searches. We find 97 white dwarf candidates with distance estimates between 10 and 120 pc, as well as 557 cool subdwarf candidates.  The subdwarfs found in this paper make up nearly half of the subdwarf systems reported from our SCR searches, and are significantly redder than those discovered thus far.  The SCR searches have now found 155 red dwarfs estimated to be within 25 pc, including 10 within 10 pc. In addition, 143 white dwarf candidates and 1155 cool subdwarf candidates have been discovered. The 1584 systems reported here augment the sample of 4724 systems previously discovered in our SCR searches, and imply that additional systems fainter than $R_{59F}$ $=$ 18.0 are yet to be discovered. ",
        "introduction": "\\label{sec:intro}  The Research Consortium On Nearby Stars (RECONS)\\footnote{\\it www.recons.org} has completed surveying the southern sky for proper motion objects with $R_{59F}$ (hereafter $R$) brighter than 16.5 mag using a specialized trawl of the individual plate detections comprising the SuperCOSMOS Sky Survey database, with new discoveries dubbed SCR (SuperCOSMOS-RECONS). These results have been detailed in six papers in the {\\it The Solar Neighborhood} (TSN) series: \\cite{2004AJ....128..437H}, \\cite{2004AJ....128.2460H}, \\cite{2005AJ....129..413S, 2005AJ....130.1658S}, \\cite{2007AJ....133.2898F}, and~\\cite{2011AJ....142...10B} (TSN VIII, X, XII, XV, XVIII, and XXV, respectively).  This seventh proper motion search paper complements the previous six by searching the entire southern sky for objects with the same lower proper motion cutoff of $\\mu$ $\\ge$ 0$\\farcs$18 yr$^{-1}$ as in previous papers, while examining objects fainter than the previous $R$ cutoff.  In this paper we stretch our search from $R =$ 16.5 to 18.0 to reveal stars intrinsically fainter than found in previous searches, as well as stars similar to previous discoveries but at larger distances. In previous searches, it was found that several stars, including SCR 1845-6357AB, the 23rd nearest system (as of 1 January 2011), had $R$ between 16.0 and 16.5, very near the search cutoff. It is natural, then, to assume that there may be more very nearby systems just beyond this cutoff that have heretofore gone unnoticed.  Despite far outnumbering their brighter counterparts, intrinsically faint nearby stars are underrepresented in the census of nearby stars, as evidenced by the discovery of many nearby red dwarfs \\citep{2006AJ....132.2360H} and white dwarfs \\citep{2009AJ....137.4547S}.  In particular, of the 256 systems with accurate parallaxes placing them within 10 pc, 174 (68\\%) have primaries of spectral type M. This search has been designed specifically to detect such faint objects. This effort is the latest in a long line of searches that revealed faint nearby stars, such as the classic sky surveys by \\cite{1971lpms.book.....G,1978LowOB...8...89G} and \\cite{1979lccs.book.....L, 1980PMMin..55....1L}.  Many more surveys have since been done utilizing modern technology but fundamentally similar techniques and are discussed in TSN XXV and references therein.  In this paper, we present photometric distance estimates for the red dwarf systems found during this search. Using the relations of \\cite{2004AJ....128..437H}, we estimate three red dwarfs to be within 25 pc, including one estimated to be at 6.7 pc. Using reduced proper motion diagrams, we also identify 97 white dwarf candidates, seven of which may be within 25 pc (see \\S~\\ref{sec:rpmd}). Particularly noteworthy for this survey, we report 557 new cool subdwarf candidates, nearly doubling the total identified during the SCR searches. The 1584 systems and 1608 objects reported in this paper bring the total number of SCR systems to 6308, consisting of 6650 objects\\footnote{Note that systems are only counted if all components are new discoveries. Multiples consisting of two SCR objects are counted once, while multiples with at least one known component are not counted towards the number of new systems.}. Identifying comprehensive samples of proper motion objects now will prepare us for large-scale astrometric surveys such as LSST and Gaia, as well as provide precursor investigation opportunities. Ultimately, these new groups of intrinsically faint objects will help develop accurate models of the Galaxy as a whole. ",
        "conclusions": "\\label{sec:discussion}   This paper completes an SCR sweep of the southern sky to $R$ $=$ 18.0 for systems with $\\mu$ $\\ge$ 0$\\farcs$18 yr$^{-1}$. In total, we have found 6308 new proper motion systems, of which 1584 are first reported in this paper, representing an increase of 33\\% over the previous count. A total of 155 red dwarfs estimated to be within 25 pc, including 10 within 10 pc, have been found, of which three and one originate from this paper, respectively. We have found an additional 97 white dwarf candidates, bringing the total to 143, and an additional 557 cool subdwarf candidates, bringing the total to 1155.  This paper is unique for the sheer numbers of white dwarf and subdwarf candidates revealed, with total counts that have tripled and doubled, respectively. Given the nature of their selection, there is obviously some contamination of the samples, but regardless, this represents a substantial increase. Of the 97 white dwarfs presented here, seven are estimated to lie within 25 pc. The subdwarfs presented here are redder and cooler than those previously discovered, helping fill in a currently underrepresented sample of cool subdwarfs discovered by these searches.  As part of our Cerro Tololo Inter-American Observatory Parallax Investigation (CTIOPI), an astrometry program carried out at the CTIO 0.9m [\\cite{{2005AJ....129.1954J}}, \\cite{2006AJ....132.2360H}, \\cite{2009AJ....137.4547S}, \\cite{2010AJ....140..897R}, and Jao et al. (2011)], we aim to obtain accurate trigonometric parallaxes for systems estimated to be near the Sun from the categories of 25 pc objects, white dwarfs, and cool subdwarfs. While the sheer number of candidates for observation means we can only select the best gems, the SCR search efforts will provide a wealth of targets for particular attention in large scale astrometric efforts such as Gaia and LSST. Such observations will help us more map the solar neighborhood more accurately than ever before."
    },
    "1107/1107.3151_arXiv.txt": {
        "abstract": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}}   We investigate the kinematic and photometric properties of the Segue 3 Milky Way companion using Keck/DEIMOS spectroscopy and Magellan/IMACS $g$ and $r$-band imaging.  Using maximum likelihood methods to analyze the photometry, we study the structure and stellar population of Segue 3.  We find the half--light radius of Segue 3 is $26''\\pm5''$ ($2.1\\pm0.4$ pc, for a distance of 17 kpc) and the absolute magnitude is a mere $M_V=0.0\\pm0.8$ mag, making Segue 3 the least luminous old stellar system known.  We find Segue 3 to be consistent with a single stellar population, with an age of $12.0^{+1.5}_{-0.4}$ Gyr and an [Fe/H] of $-1.7^{+0.07}_{-0.27}$.  Line--of--sight velocities from the spectra are combined with the photometry to determine a sample of 32 stars which are likely associated with Segue 3.  The member stars within three half--light radii have a velocity dispersion of $1.2\\pm2.6$ \\kms.  Photometry of the members indicates the stellar population has a spread in [Fe/H] of $\\lesssim0.3$ dex.  These facts, together with the small physical size of Segue 3, imply the object is likely an old, faint stellar cluster which contains no significant dark matter.  We find tentative evidence for stellar mass loss in Segue 3 through the eleven candidate member stars outside of three half-light radii, as expected from dynamical arguments.  Interpretation of the data outside of three half--light radii, is complicated by the object's spatial coincidence with a previously known halo substructure, which may enhance contamination of our member sample. ",
        "introduction": "\\label{sec:intro}    Segue 3 (hereafter Seg\\,3) is a recently discovered Milky Way satellite, initially estimated to have an extremely low luminosity ($M_{\\rm V} \\sim -1.2$, $\\sim$ 250 $L_{\\Sun}$), a half-light radius of $\\sim$ 3 pc, and a distance of $\\sim$17 kpc \\citep{belokurov10a}.  Five other Milky Way satellites are known to have comparably small luminosities and scale sizes: Whiting 1 ($M_{\\rm V} = -2.46$), Pal 1 ($M_{\\rm V} = -2.52$), AM 4 ($M_{\\rm V} \\sim -1.8$), Koposov 1 ($M_{\\rm V} = -1.35$), and Koposov 2 ($M_{\\rm V} \\sim -0.35$) \\citep[][2010 edition]{harrisGCcat}. All six of these objects are relative outliers in size and luminosity from other known Milky Way satellites, with luminosities a factor of three lower than those of the other 112 Milky Way satellites with half-light radii smaller than 5 pc.  These six extreme satellites lie at distances ranging from 17 kpc (Segue 3, Pal 1) to 40 kpc or more (Koposov 1 and 2).  Despite these halo distances, one explanation for their anomalously low luminosities is stellar mass loss owing to dynamical evolution, such as tidal stripping or tidal shocking.  At their current stellar masses and sizes, the evaporation timescales of these clusters are also shorter than the observed ages of the systems (see e.g. discussion in \\citealt{koposov07b}), suggesting that evaporation is playing a major role in their evolution.  The hypothesis that these satellites are experiencing substantial stellar mass loss is supported by observations: Strong evidence for extra-tidal stars has been found in Pal 1, AM 4, and Whiting 1 \\citep{ carraro09a, carraro07a,NO10a}. (Koposov 1 and 2 haven't yet been studied thoroughly enough to confirm or rule out the presence of extra-tidal stars.)  Most of these six ultra-low luminosity satellites have been (at least tentatively) associated with larger scale stellar streams in the halo. For example, Whiting 1 and Koposov 1 have been associated with the Sgr stream \\citep{carraro07a, koposov07b}, AM 4 possibly with the Sgr stream \\citep{carraro09a}, Pal 1 possibly with the Galactic Anticenter Stellar Structure \\citep[GASS; ][]{frinchaboy04} or Canis Major \\citep{forbes10a}, and Segue 3 possibly with the Hercules-Aquila cloud \\citep{belokurov10a}.   Perhaps all six of these $M_V > -2.5, r_{1/2} < 5$ pc satellites are stripped down versions of more luminous objects, and can provide insight into the evolution of the Milky Way's satellite population and into the build-up of the halo. Because of their small physical sizes and luminosities alone, they have all been classified as globular clusters in the literature.  Their central surface brightnesses are higher than those of the known Milky Way dwarfs, so they are not simply stripped versions of those objects.  Although this circumstantial evidence supports a model where these outliers may instead be stripped versions of more luminous globular clusters, no robust classification or detailed kinematic study of these objects has yet been done.  We aim to improve our understanding of these unusual Milky Way satellites by conducting a detailed kinematic study of Segue 3, the first to be done for any comparable object. In particular, we study Magellan/IMACS photometry in $g$ and $r$ of Segue 3 to investigate its structural parameters and star formation history.  We combine this photometry with Keck/DEIMOS spectroscopy to obtain a robust member sample.  We use these spectroscopic members to study Seg 3's internal kinematics and to look for evidence of the extra-tidal stars seen in similar clusters. We use its internal kinematics, improved structural parameters, and star formation history to infer a star cluster classification for Seg\\,3. ",
        "conclusions": "We have investigated the nature of the faint Milky Way satellite Segue 3 using new Keck/DEIMOS spectrocopic and Magellan/IMACS photometric data.  Using the data we have identified a sample of 32 probable member stars for Seg\\,3, one of which is definitively identified as a binary system through repeated velocity measurements.  We summarize our conclusions as follows:  \\begin{itemize}  \\item New IMACS photometry reveals a distinct main sequence of stars in the CMD of Seg\\,3, with the majority of these stars lying within two half--light radii of the center.  Using maximum likelihood methods which fit isochrones from \\citet{dotter08}, we infer an average age of $12.0^{+1.5}_{-0.4}$\\,Gyr and an average [Fe/H] of $-1.7^{+0.07}_{-0.27}$ for the stellar population.  \\item We analyze the structural properties of Seg\\,3, using Plummer and exponential profiles.  We find structural parameters similar to those of \\citet{belokurov10a}, except for the half-light radius which we find to be smaller by $\\sim 30\\%$.  Using a Salpeter IMF and CMD properties from our photometry, we find Seg\\,3 has an absolute magnitude of $M_{V}=0.0\\pm0.8$, making it the faintest stellar system known to date.  \\item Examining 20 spectroscopic member stars within three half--light radii, we find a kinematically bound group of 19 member stars and identify one star which is an outlier of the velocity distribution (at 95\\% confidence).  For this group of bound stars we infer a velocity dispersion of $1.2\\pm2.6$ \\kms, consistent with very low values for the true velocity dispersion.  \\item Segue 3 is likely a faint stellar cluster with no significant dark matter content.  This is evidenced by our measurement of the velocity dispersion within three half--light radii, the (relatively) high average and low scatter in [Fe/H] values implied by isochrones, and small physical size of Seg\\,3 relative to similar luminosity ultra--faint dwarf galaxies.  \\item We find signs of possible mass loss in Segue 3, indicated by eleven candidate member stars outside of $3r_{1/2}$. While these stars may originate from the Milky Way or nearby substructure, we estimate a Jacobi radius of $r_{\\rm J}<5r_{1/2}$ and a half--mass relaxation time of $t_{\\rm rh}\\sim150$ Myr for Seg\\,3, suggesting mass loss should be important.  \\item Seg\\,3 appears spatially coincident with an extended overdensity discovered by \\citet{dejong10a}, which may be associated with the Hercules--Aquila cloud.  However, we find it is kinematically offset from the structure; evidenced by the low density of non-member stars within $\\pm45$\\kms\\ from the systemic velocity of Seg\\,3.  In our spectroscopic sample we identify a collection of stars with line--of--sight velocities $<200$\\kms, which may be associated with the structure.  We find at most one star from this population may be misidentified as a member.  \\end{itemize} \\phantom{need space between conclusions and acknowledgments}"
    },
    "1107/1107.5059.txt": {
        "abstract": "We propose a simple model for the origin of fast and slow rotator early-type galaxies (ETG) within the hierarchical $\\Lambda$CDM scenario, that is based on the assumption that the mass fraction of stellar discs in ETGs is a proxy for the specific angular momentum expressed via $\\lambda_R$.  Within our model we reproduce the fraction of fast and slow rotators as a function of magnitude in the \\atl survey, assuming that fast rotating ETGs have at least $10\\%$ of their total stellar mass in a disc component. In agreement with \\atl observations we find that slow rotators are predominantly galaxies with $ M_* > 10^{10.5}$ M$_{\\odot}$ contributing $\\sim 20 \\%$ to the overall ETG population.  We show in detail that the growth histories of fast and slow rotators are  different, supporting the classification of ETGs into these two categories. Slow rotators accrete between $\\sim 50 \\% -90\\%$ of their  stellar mass from  satellites and their most massive progenitors have on average up to 3 major mergers during their evolution. Fast rotators in contrast, accrete less than $50\\%$ and have on average less than one major merger in their past.  We find that the underlying physical reason for the different growth histories is the slowing down and ultimately complete shut-down of gas cooling in massive galaxies. Once cooling and associated star formation in disc stops, galaxies grow via infall from satellites. Frequent minor mergers thereby, destroy existing stellar discs via violent relaxation and also tend to lower the specific angular momentum of the main stellar body, lowering $\\lambda_R$ into the slow rotator regime.  On average the last gas-rich major merger interaction in slow rotators happens at $z > 1.5$, followed by a series of minor mergers. These results support the idea that kinematically decoupled cores (KDC) form during gas-rich major mergers at high-z followed by minor mergers, which build-up the outer layers of the remnant, and make remnants that are initially too flat compared to observations become rounder. Fast rotators are less likely to form such KDCs due to the fact that they have on average less than one major merger in their past.  Fast rotators in our model have different formation paths. The majority, $78 \\%$  has bulge-to-total stellar mass ratios $B/T> 0.5$ and managed to grow stellar discs due to continued gas cooling or bulges due to frequent minor mergers. The remaining $22 \\%$ live in high density environments and consist of low $B/T$ galaxies with gas fractions below $15 \\%$, that have exhausted their cold gas reservoir and have no hot halo from which gas can cool. These fast rotators most likely resemble the flattened disc-like fast rotators in the \\atl survey.  Our results predict that ETGs can change their state from fast to slow rotator and vice versa, while the former is taking place predominantly at low z $(z < 2)$, the latter is occurring during cosmic epochs when cooling times are short and galaxies gas-rich. We predict that the ratio of the number density of slow to fast rotators is a strong function of redshift, with  massive $( > 10^{10}$ M$_{\\odot}$) fast rotators being more than one order of magnitude more frequent at $z \\sim 2$. ",
        "introduction": "Much attention has been paid to the modelling of the formation of early-type galaxies (ETG) over the last three decades \\citep[e.g.][]{1972ApJ...178..623T,1978MNRAS.184..185W,1983MNRAS.205.1009N,1992ARA&A..30..705B,1996ApJ...464..641M,2003ApJ...597..893N,2004A&A...418L..27B,2006ApJ...650..791C,2010ApJ...723..818H}. Once believed to be coeval featureless stellar systems in virial equilibrium that formed from a 'monolithic' collapse \\citep{1974MNRAS.166..585L}, an  ever-increasing wealth of observational data has revealed a multitude of distinct physical properties: in their rotational support \\citep{1978MNRAS.183..501B,1983ApJ...266...41D}, isophotal shape \\citep{1989A&A...217...35B,1996ApJ...464L.119K}, photometric profiles \\citep{1994AJ....108.1598F,1995AJ....110.2622L,2009ApJS..182..216K}, and in their stellar angular momentum \\citep{2007MNRAS.379..401E,2007MNRAS.379..418C}.  Driven by such observations theoretical models have been put to the test, and in particular the $\\Lambda$CDM paradigm for structure formation itself  within which structure grows hierarchically via a sequence of mergers and accretion events \\citep{1978MNRAS.183..341W}.  Early numerical N-body simulations \\citep{1972ApJ...178..623T} identified galaxy mergers as a viable mechanism to transform dynamical cold stellar systems into dynamical hot ones resembling early-type galaxies with a de Vaucouleurs-like  light profile \\citep[see also][]{1992ARA&A..30..705B,2006MNRAS.369..625N}. The main driver for this morphological transformation being violent relaxation \\citep{1967MNRAS.136..101L} during which stars get redistributed. The merger scenario subsequently has been tested against two main categories of observational data: global properties in large homogenous surveys \\citep[e.g.][]{2003AJ....125.1817B}  and detailed high-resolution observations of an inhomogeneous set of individual galaxies \\citep[e.g.][]{2002MNRAS.329..513D}.  Galaxy formation models, in particular semi-analytic models,  employing the merger scenario, reproduce successfully the mass function, colour distribution, star formation history  and metallicity of early-type galaxies as observed e.g. in the Sloan Digital Sky Survey (SDSS) \\citep[e.g.]{2006MNRAS.366..499D}. However, they still have problems in e.g. recovering the trend in $\\alpha/Fe$ with stellar velocity dispersion \\citep{2005ApJ...621..673T,2005MNRAS.363L..31N}  or the joint distribution in sizes and velocity dispersion at $z=0$ \\citep{2009ApJ...698.1232V,2010MNRAS.405..948S}. The latter has also been reflected in mismatches to the fundamental plane  of early type galaxies \\citep{2007MNRAS.376.1711A}. While the deviations regarding the sizes and velocity dispersions likely stem from too efficient in-situ star formation in progenitors of present-day early-type galaxies \\citep{2010MNRAS.405..948S}, the failure to produce enough $\\alpha$-elements is most likely  connected to a combination of too much residual star formation \\citep{2005ApJ...621..673T} and satellite infall \\citep{ks06}.  The inner parts of early-types show further detailed characteristics in the form of steep increasing  power-law or  core like profiles \\citep{1994AJ....108.1598F,1995AJ....110.2622L,2009ApJS..182..216K}. The former has been associated with cold gas in discs that loses its angular momentum during the merging process \\citep[e.g.][]{1991ApJ...370L..65B}  and settles to the centre of the remnant where it triggers a star burst, thereby steepening the light profile \\citep{2008ApJ...679..156H}.  Core galaxies on the other hand are most likely the product of a {\\it non-dissipative} merger event between two galaxies hosting super-massive black holes (SMBH) at their centre \\citep{2001ApJ...563...34M,2005MNRAS.359.1379K,2009ApJS..181..486H}. During the inward spiralling the SMBH pair kick out stars via three-body interaction and produces a density core in the centre of the remnant \\citep{2001ApJ...563...34M}.  Classifying elliptical galaxies by the deviation of their isophotal shape from pure ellipses has been suggested as a method to separate the population of elliptical galaxies into two classes with box-like (boxy) and disk-like deviations (disky) \\citep{1989A&A...217...35B}. Numerical simulations show that equal mass mergers in general produce remnants with boxy isophotes due to  stronger gravitational interaction and associated violent relaxation \\citep{2000MNRAS.316..315B,2003ApJ...597..893N} efficiently destroying progenitor discs. In contrast unequal mass major mergers  with $2 \\le M_1/M_2 \\le 4$ result on average in disky remnants for the majority of projections under which they are viewed \\citep{2003ApJ...597..893N}. Observationally, the average isophotal shape of elliptical galaxies turns toward boxy going up in mass \\citep[see howerver][]{pap3} indicating a transition in the formation process of massive ellipticals that cannot be attributed to the mass ratio during the last major merger alone, which in general does not change significantly enough in frequency as a function of mass  \\citep{2005MNRAS.359.1379K}. In addition unequal mass mergers of early-type galaxies are able to transform disky ellipticals into boxy ones in the majority of simulated cases \\citep{2006ApJ...636L..81N}. Within a cosmological context the relative fraction of dry, early-type major mergers increases with stellar mass of the remnant \\citep{2003ApJ...597L.117K} providing the additional channel needed to recover the observed trend with mass \\citep{2005MNRAS.359.1379K}. In a detailed study of the orbital content of stars in merger remnants \\citet{2005MNRAS.360.1185J} showed that the individual mixture of orbits is the deciding factor on the isophotal shape of disky and boxy elliptical galaxies.  Making full use of the dynamical information available in integral-field observations the SAURON sample \\citep{2002MNRAS.329..513D} measured the specific baryonic angular momentum content via the proxy $\\lambda_R$ within one $R_e$ of a representative sample of early-type galaxies and separate them into fast and slow rotators according to their $\\lambda_R$ value, in a robust way that is nearly insensitive to projection effects  \\citep{2007MNRAS.379..401E,2007MNRAS.379..418C}. Observationally, fast rotators present regular rotation patterns aligned with the photometry, while slow rotators have low angular momentum content and show misalignments between the photometry and the velocity axes and often exhibit kinematically distinct cores \\citep{2008MNRAS.390...93K}. The division into fast and slow rotators is further supported by recent simulations of \\citep[hereafter Paper VI]{pap6} that show that fast (resp. slow) rotators have a high (resp. low) angular momentum content, their photometric and kinemetric position angle are aligned (resp. misaligned), they present (resp. do not present) regular velocity patterns. While the SAURON survey was biased towards massive ETGs \\citep{2002MNRAS.329..513D} the \\atl survey \\citep[][hereafter Paper I]{pap1} for the first time allows to give a complete view of the local ETG population within a radius of $42 $ Mpc using the full information of integral-field observations. The combination of completeness with detailed structural information and complementary observations on e.g. the HI \\citep{paolo} or molecular  \\citep{pap4}  gas masses, make the \\atl survey an ideal data set to investigate the formation of ETGs.  In this paper we attempt to model fast and slow rotators within the standard $\\Lambda$CDM paradigm and analyse specific differences in their formation. The structure of the paper is as follows: we start by summarising the basic model ingredients used in our semi-analytic modelling (SAM) approach  in section \\ref{mod}, followed by showing the performance of the model in section \\ref{modnor}. In section \\ref{fas} we introduce our model for fast and slow rotators and investigate the robustness of our assumptions before we analyse in detail differences in their formation in section \\ref{demo} and predict their redshift evolution in section \\ref{z}.  Section \\ref{var} is devoted to the question of cosmic variance in the \\atl sample followed by sections \\ref{sum} \\& \\ref{con} in which we summarise our results and conclude respectively. ",
        "conclusions": "\\label{con} In this paper we made use of the completeness of the \\atl sample to investigate the origin and formation history of fast and slow rotator ETGs within a self-consistent cosmological framework using the semi-analytical modelling approach. We  here present a model in which  the difference between fast and slow rotator ETGs is purely based on the stellar disc fraction found in them, and predict that fast rotators have disc fractions  $M_{disc}/M_* > 0.1$. We find that  slow rotators within an evolving universe mark the transition in the ability of galaxies to cool  gas and to rebuild stellar discs.  We find a clear separation in the growth history of fast and slow rotator ETGs, supporting the observationally motivated distinction. In particular the accreted fraction of stars shows a clear distinction between fast and slow rotators, with the latter having between $50 \\% - 90 \\%$ of their stellar mass  accreted from satellites while the former has $ < 50 \\%$ accreted.  Although we find a clear separation into fast and slow rotators, we also find that ETGs can switch their state between fast and slow rotator and vice versa  based on stellar disc growth or destruction by mergers, suggesting that fast and slow rotator ETGs are transient. These changes however, occur predominantly at higher redshifts when cooling is more efficient and mergers more frequent. The fraction of slow rotators,  shows a continued increases with time due to the conversion of fast rotators and the inability of gas to cool and convert slow rotators back to fast rotators. Massive present-day slow rotators therefore, can be viewed  as the final stage in the evolution of  ETGs.  Future high redshift observations of ETGs will be able to reveal any possible evolution in the fraction of fast and slow rotators and test the presented model  further."
    },
    "1107/1107.1224_arXiv.txt": {
        "abstract": "We present observations of the intermediate to massive star-forming region I05345+3157 using the molecular line tracer CS($2-1$) with CARMA to reveal the properties of the dense gas cores. Seven gas cores are identified in the integrated intensity map of CS($2-1$).  Among these, core 1 and core 3 have counterparts in the $\\lambda$~=~2.7~mm continuum data. We suggest that core 1 and core 3 are star-forming cores that may already or will very soon harbor young massive protostars. The total masses of core 1 estimated from the LTE method and dust emission by assuming a gas-to-dust ratio are 5$\\pm$1 M$_{\\sun}$ and 18$\\pm$6 M$_{\\sun}$, and that of core 3 are 15$\\pm$7 M$_{\\sun}$ and 11$\\pm$3 M$_{\\sun}$.  % The spectrum of core 3 shows blue-skewed self-absorption, which suggests gas infall -- a collapsing core. The observed broad linewidths of the seven gas cores indicate non-thermal motions. These non-thermal motions can be interactions with nearby outflows or due to the initial turbulence; the former is observed, while the role of initial turbulence is less certain. Finally, the virial masses of the gas cores are larger than the LTE masses, which for a bound core implies a requirement on the external pressure of $\\sim 10^{8}$ K cm$^{-3}$. The cores have the potential to further form massive stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5820_arXiv.txt": {
        "abstract": "We present a novel and flexible tensor approach to computing the effect of a time-dependent tidal field acting on a stellar system. The tidal forces are recovered from the tensor by polynomial interpolation in time. The method has been implemented in a direct-summation stellar dynamics integrator (\\nbody) and test-proved through a set of reference calculations: heating, dissolution time and structural evolution of model star clusters are all recovered accurately. The tensor method is applicable to arbitrary configurations, including the important situation where the background potential is a strong function of time. This opens up new perspectives in stellar population studies reaching to the formation epoch of the host galaxy or galaxy cluster, as well as for star-burst events taking place during the merger of large galaxies. A pilot application to a star cluster in the merging galaxies NGC~4038/39 (the Antennae) is presented. ",
        "introduction": "The halos of nearly all galaxies are populated by old globular clusters that presumably formed in gaseous discs at high redshift \\citep[$z\\sim3-5$,][] {2005ApJ...623..650K}. Young `populous clusters', or `super-star clusters' are found in the Large Magellanic Cloud \\citep{1961ApJ...133..413H}, starburst galaxies \\citep[e.g.][]{1971A&A....12..474V}, interacting galaxies and merger remnants \\citep[e.g.][]{1995AJ....109..960W,1997AJ....114.2381M} and also in quiescent spirals \\citep{2002AJ....124.1393L, 2008IAUS..250..247F}. This suggests that globular cluster formation is not unique to the early Universe and that the formation of these dense stellar systems is a common phenomenon in star formation \\citep*{1997ApJ...480..235E, 2010ARA&A..48..431P}.  There is increasing evidence from studies of the Milky Way and the Andromeda galaxy that some globular clusters have only recently (past $\\sim\\,$Gyr) been brought in through satellite accretion \\citep[e.g.][]{2009ApJ...694.1498M,2010ApJ...717L..11M}. In order to understand the relation between the young massive clusters and the old globular clusters, i.e. their life-cycle, we need to place their evolution in a cosmological context. During the formation process of galaxies through repeated accretion phases, substructures such as star clusters or dwarf satellites evolve along complex orbits in a non-static external potential. This makes their evolution difficult (if not impossible) to describe analytically.  To approach this problem numerically is also challenging because of the large range of evolutionary time scales involved, ranging from several days for tight binaries in the cores of clusters to a Hubble time for the host galaxy \\citep[see the recent review by][]{2011EPJP..126...55D}. To be able to self-consistently model the evolution of star clusters in `live' galaxies one needs to rely on a direct-tree hybrid approach  \\citep[e.g.][]{2007PASJ...59.1095F}.  This is why most studies of the (dynamical) evolution of star clusters simplify the effect of the external tidal field by assuming a static background potential \\citep[e.g.][]{1990ApJ...351..121C, 1997MNRAS.289..898V,2001ApJ...561..751F, 2003MNRAS.340..227B,2004AJ....127.2753D,2008MNRAS.389L..28G,2008AJ....135.2129H,2009MNRAS.399.1275P,2011MNRAS.411.1989Z}. Although this is probably an adequate approximation for many purposes, it does not suffice for more complicated orbits such as those of clusters in mergers of massive galaxies, or satellite galaxies that are accreted.  In light of the last point, several recent studies have adopted a (semi-)analytical approach to star cluster evolution in more realistic external tides, such as the hierarchical build-up of galaxies \\citep[e.g.][]{2008ApJ...689..919P} and galaxy mergers \\citep[e.g.][]{2011MNRAS.414.1339K}. In here the effect of mass loss because of stellar evolution, evaporation of stars over the tidal radius and the shock-enhanced escape of stars because of rapidly varying tidal fields (i.e. disc crossings and bulge shocks) are applied analytically. In almost all cases these processes are assumed to be independent of each other such that the individual resulting mass-loss rates are simply added.  An important, and often dominant, mass-loss mechanism is the relaxation driven escape of stars, so-called evaporation. Some models assume that a constant fraction of the stars evaporate per half-mass relaxation time, independent of the orbit \\citep[e.g.][]{1997ApJ...474..223G,2008ApJ...689..919P}. Others consider that the escape fraction depends on the galactocentric radius, often assumed to be the pericentre distance \\citep{1962AJ.....67..471K,1983AJ.....88..338I,2001ApJ...561..751F}. The final lifetime of clusters can be quite different, depending on the details of the assumptions that are made.  Another critical disruptive agent is mass-loss  due to external tidal perturbations, or `shocks'. The related lifetime scales linearly with the density of the cluster \\citep{1958ApJ...127...17S,1972ApJ...176L..51O} and it is, therefore, that the results of semi-analytical models rely critically on what is assumed for the relation between the cluster mass and the half-mass radius; a relation which is not only affected by evaporation (i.e. because of the reduction of the mass in time), but also because of relaxation driven expansion \\citep{1965AnAp...28...62H, 1984ApJ...280..298G,2002MNRAS.336.1069B,2010MNRAS.408L..16G}, which as far as we are aware is neglected in all semi-analytic approaches. For clusters on circular orbits in static potentials that are well within the tidal limit initially (the half-mass radius being less than a few percent of the tidal radius), it was found that the expansion phase dominates the evolution in the first half of the cluster's lifetime, while evaporation dominates in the second half \\citep{2011MNRAS.413.2509G}. In the latter stage the cluster density adjusts to the mean (galactic) density along the orbit \\citep{2010MNRAS.407.2241K}. It is, therefore, necessary to consider both the internal evolution (relaxation) and the external effects (tides) simultaneously if one considers the entire life-cycle of clusters.  Current direct $N$-body codes are capable of solving the $N$-body problem numerically for $N\\sim10^5$, together with the effects of mass-loss of the individual stars, binary interactions and tidal fields \\citep{2001MNRAS.321..199P,2003gnbs.book.....A}. Thanks to increasing computational power, it is now possible to combine the relaxation driven evolution of star clusters in cosmologically motivated external conditions. In this paper we present a new method that integrates the effect of any tidal field to the evolution of galaxy substructures, that we specifically apply to star clusters. We do this by extracting the tidal tensor, that contains all information about the tidal field at the location of the substructure, from galaxy simulations and subsequently `feed' this to a stellar dynamics code. However, we have not yet implemented the other half of the scale coupling, as the feedback from the small scales (e.g. metal enrichment, stellar winds, supernova explosions) is not retroceded to the ambient galactic medium.  The paper is organised as follows: first, we setup the framework for the computation of the tidal acceleration by means of tensors (Section~\\ref{sec:generalcase}). The expressions found are then applied to the special cases of circular orbits: well-known expressions are retrieved using the new formalism in Section~\\ref{sec:circular}. The role of the galactic profile on the evolution of clusters through the escape of stars is specially explored in Section~\\ref{sec:escape}. The limitations of the analytical approach are explained in Section~\\ref{sec:noncircular}, while Section~\\ref{sec:nbody6tt} presents the numerical implementation of the method to compute the tidal forces in a stellar dynamics code. A comparison with previous results obtained for idealized configurations is carried out. Applications to innovative cases are presented as the first practical illustrations of the method. Finally, the limitations and some possible future developments of our approach are discussed. ",
        "conclusions": "In this contribution, we propose a new formalism to describe the tidal field in $N$-body simulations by the means of tensors. Although the analytical approach rapidly becomes too involved for accelerated motions, we have derived the expressions of quantities representing the effect of the tides on stellar systems for circular orbits, with no restriction on the shape of the external potential. The main results of this study are: \\begin{itemize} \\item The use of the tidal and effective tensors allows us to simplify the representation of the problem, without loss of information (Section~\\ref{sec:tensors}). Useful quantities like the tidal radius (equation~\\ref{eqn:rt}), the energy of the Jacobi surface (equation~\\ref{eqn:ej}) or the escape timescale (equation~\\ref{eqn:escapetime}) can be easily computed for a given cluster. \\item The tidal radius is a first order approximation of the effect of a galaxy on a star cluster. The three-dimensional shape of the galactic potential has a second order effect (Section~\\ref{sec:escape} and equation~\\ref{eqn:dissolution}). \\item For a given tidal radius, a point-mass is the most efficient galactic profile in dissolving a cluster. Shallower potentials lead to longer dissolution times (\\fig{dissolutiontime}). \\item A cluster in compressive tidal mode loses its stars more rapidly than in isolation, but significantly more slowly than if it was in extensive mode (Section~\\ref{sec:escape}). \\item The knowledge of the evolution of the cluster parameters (half-mass radius, mass, energy distribution) is key to estimate the cluster life-time (Section~\\ref{sec:escape} and Appendix~\\ref{app:flux}). \\item For non-circular orbits, the analytical approach becomes very involved so that general and simple formulae do not exist (Section~\\ref{sec:noncircular}). \\end{itemize} To overcome this issue, we have developed and implemented a numerical method called \\nbodytt which computes the evolution of $N$-body models of star clusters in any tidal field, by means of the pre-calculation of inertial tidal tensors. This method has been successfully tested and applied to innovative cases. Our main conclusions are: \\begin{itemize} \\item The tidal force felt by the stars of a star cluster can be accurately computed by evaluating the tidal tensor at all time by means of quadratic interpolation (Section~\\ref{sec:tests}). \\item The dependence of the dissolution time on the three-dimensional shape of the galactic potential highlighted in Section~\\ref{sec:escape} is confirmed by our numerical experiments. The numerical method does not suffer from the lack of knowledge of the evolution of the cluster properties that one has to face in a (semi-)analytical approach (Section~\\ref{sec:otherprofiles}). \\item Our implementation has been applied to the complex case of a cluster in the major merger of the Antennae galaxies (Section~\\ref{sec:antennae}). Its mass-loss reflects the nature of the time-dependent tidal field experienced along the orbit. In particular, the alternation of extensive and compressive tidal modes strongly affects the instantaneous dissolution rate, over a time-scale of several $10^7 \\yr$. \\end{itemize}  This preliminary study shows the way to a very wide range of possible applications. However, one should keep in mind that our study is limited in several aspects. On the one hand, our cluster simulations are gas-free. Therefore, we cannot address the important point of the early life of the star cluster, prior to gas expulsion (first $\\sim 10^6 \\yr$ of the cluster's life), and we limit our study to initially relaxed systems. On the other hand, the second half of the galaxy-cluster coupling is not taken into account by our method: the feedback from stellar evolution, the escape of stars in the field and the formation of long tidal tails or streams are not implemented. Although they have a limited impact on the evolution of the cluster itself, it would be important to monitor their effects at larger scale (e.g. the metal enrichment of the interstellar medium). Furthermore, the mass-loss experienced by a cluster could affect its orbit within its host galaxy and thus, change the tidal field. In our approach, the scale decoupling does not allow to account for such effect.  To conclude, we have shown that a star cluster plunged in a time-dependent tidal field leads to a complex evolution, out-of-reach of (semi-)analytical approaches. A numerical method like the one proposed by \\nbodytt provides a framework for future explorations of the role of the tides on star clusters. Among others, the use of a stellar mass function, primordial binaries and high order stars, stellar evolution and stellar mass-loss are as many lines of investigation to be pursued on top of the evolution in a background, time-dependent, galactic potential.  In a forthcoming paper, we will use more detailed galactic simulations including a prescription on the star formation, to compute the tidal tensors of an entire cluster population. \\nbodytt will help us to derive the cluster- mass and age functions and their evolution, in various type of galaxies, in particular in mergers."
    },
    "1107/1107.2777_arXiv.txt": {
        "abstract": "We have recently proved the impossibility of imposing the condition of local charge neutrality in a  self-gravitating system of degenerate neutrons, protons and electrons in $\\beta$-equilibrium. The coupled system of the general relativistic Thomas-Fermi equations and the Einstein-Maxwell equations have been shown to supersede the traditional Tolman-Oppenheimer-Volkoff equations. Here we present the Newtonian limit of the new equilibrium equations. We also extend the treatment to the case of finite temperatures and finally we give the explicit demonstration of the constancy of the Klein potentials in the case of finite temperatures generalizing the condition of constancy of the general relativistic Fermi energies in the case of zero temperatures. ",
        "introduction": "\\label{sec:1}  It is well known that the classic work of Oppenheimer and Volkoff \\cite{oppenheimer39} addresses the problem of neutron star equilibrium configurations composed only of neutrons. For the more general case when protons and electrons are also considered, in nearly all of the scientific literature on neutron stars it is assumed that the condition of local charge neutrality applies identically to all points of the equilibrium configuration (see e.g.~\\cite{nsbook}). Consequently, the corresponding solutions of the Einstein equations for a non-rotating neutron star, following the work of Tolman \\cite{tolman39} and of Oppenheimer and Volkoff \\cite{oppenheimer39}, have been systematically applied. However, the necessity of processes leading to electrodynamical phenomena during the gravitational collapse to a black hole \\cite{physrep} suggests a critical reexamination of the current treatment of neutron stars. In this regard in a set of recent articles (see e.g.~\\cite{PRC2011,PLB2011}) we have developed the first steps toward a new consistent treatment for the description of neutron stars overcoming the traditional Tolman-Oppenheimer-Volkoff equations.  We have recently proved the impossibility of imposing the condition of local charge neutrality in a self-gravitating system of degenerate neutrons, protons and electrons in $\\beta$-equilibrium \\cite{PLB2011}. It has been there discussed that the traditional approach for the description of neutron stars adopting the condition of local charge neutrality is inconsistent with the Einstein-Maxwell equations and with micro-physical conditions of equilibrium within the quantum statistics. We emphasize the basic role of the constancy of the general relativistic Fermi energy of each species. Consequently, the traditional Tolman-Oppenheimer-Volkoff system of equations has been superseded by the coupled system of the general relativistic Thomas-Fermi equations and the Einstein-Maxwell equations. We have been guided in this new approach by the  generalization of the Feynman-Metropolis-Teller treatment of compressed atoms to the relativistic regimes \\cite{PRC2011}. There, using the existence of scaling laws \\cite{ruffini2008,popov2010,popov2011}, the results were extended from heavy nuclei to the case of globally neutral nuclear matter cores of stellar dimensions with mass numbers $A\\approx (m_{Planck}/m_n)^3\\approx10^{57}$ or $M_{core}\\approx M_{\\odot}$.  In the present article we first recall in Sec.~\\ref{sec:2} the new system of general relativistic  Einstein-Maxwell-Thomas-Fermi governing a self-gravitating system of neutrons, protons and electrons in $\\beta$-equilibrium. We then address three new aspects of the problem: 1) we study the Newtonian limit of the equilibrium equations (see Sec. III); 2) we consider the proper generalization of the above treatment to the case of finite temperatures (see Sec. IV); 3) we generalize and prove the constancy of the Klein potentials of each species generalizing the condition of constancy of the general relativistic Fermi energies encountered in the case of zero temperatures (see Sec.~\\ref{sec:4}). We finally show that the thermal effects do not affect the electrodynamical structure of the equilibrium configurations and the results are qualitatively and quantitatively quite similar to the ones obtained with the degenerate approximation (see Sec.~\\ref{sec:4}). ",
        "conclusions": "\\label{sec:5}  In this article we have addressed three new aspects of the description of a self gravitating system of neutrons, protons and electrons in $\\beta$-equilibrium:  1) We have presented the Newtonian limit of the treatment by taking the weak field approximation and the non-relativistic $c \\to \\infty$ limit of the general relativistic Thomas-Fermi and Einstein-Maxwell equations (\\ref{eq:Gab12})--(\\ref{eq:GRTF2}). The numerical integration of the Newtonian equations shows that the gravito-electrodynamic structure evidenced in \\cite{PLB2011} (see also Sec.~\\ref{sec:2}) is not of general relativistic nature but is already present in the Newtonian regime. However, in view of the large quantitative discrepancies of the Newtonian regime in the description of other neutron star properties like mass and radius (see e.g.~Fig.~3 in \\cite{merloni98}), a general relativistic treatment is mandatory in any astrophysical consideration.  2) It has been also presented the extension of the previous treatment \\cite{PLB2011} to finite temperatures. Although the thermal energy expected to be stored in old neutron stars with surface temperatures $\\sim 10^6$ K \\cite{crab1,crab2} is much larger than the internal Coulomb energy (see Sec.~\\ref{sec:4}), still the electromagnetic structure (see Fig.~\\ref{fig:EGRNewtonian}) is unaffected by the presence of the thermal component. Physically this is due to the very large Fermi energy of the neutrons $\\sim 1$ GeV, of the protons $\\sim 10$ MeV and of the electrons $\\sim 0.1$ GeV, as can be seen from Eq.~(\\ref{eq:Tfermi}).  3) More generally, we have given the explicit demonstration of the constancy throughout the configuration of the Klein potentials of each species present in the system in the more general case of finite temperatures. This generalizes the condition of the constancy of the general relativistic Fermi energies derived in the special case $T=0$ in \\cite{PLB2011}.  All these results are relevant to the generalization of the Feynman-Metropolis-Teller treatment of compressed atoms to relativistic regimes in presence of thermal effects and they will necessarily also apply in the case of additional particle species and of the inclusion of strong interactions in neutron stars \\cite{PLB3011}."
    },
    "1107/1107.4151_arXiv.txt": {
        "abstract": " ",
        "introduction": "} The \\c supernova remnant (SNR) (G119.5+10.2) is a composite \\s characterized by a large radio shell enclosing a smaller pulsar wind nebula (PWN).  \\cite{Pineault1993} derived a kinematic distance of $1.4\\pm0.3$ kpc based on associating an \\ion{H}{1} shell found to the northwestern part of the remnant to the remnant itself. Observations in X-rays with \\textit{ASCA} and \\textit{ROSAT} revealed a central filled SNR with emission extending to the radio shell in the south and southeast as well as in the north and northwest radio-quiet regions \\citep{Seward1995}. X-ray observations with {\\em ROSAT} also revealed the point source RX J0007.0+7302 \\citep{Seward1995}. X-ray observations with \\C revealed a compact PWN and a jet-like structure \\citep{Halpern2004}. Radio and X-ray characteristics of \\c imply an age in the range 5000--15000 years \\citep{Pineault1993,Slane1997,Slane2004}. \\cite{Slane2004} estimated for the age of the SNR a value of $1.3 \\times 10^4 \\, d_{1.4} $ yr (where $d_{1.4}$ is the distance in units of 1.4 kpc) which is in good agreement with the spin-down age estimate of 14,000 yr from \\g-rays \\citep{CTA1}.  The offset of the point source \\r from the geometrical center of the radio SNR allows for the estimate of the transverse velocity of the point source which \\cite{Slane2004} estimated to be $\\sim$450 km s$^{-1}$.  Prior to the launch of \\fermis the EGRET \\g-ray source 3EG J0010+7309, which lies within the boundaries of the radio SNR, showed the characteristics of a pulsar \\citep{Brazier1998,Halpern2004}. \\cite{mattox1996_pulsar} discussed this source as a potential candidate for a radio-quiet \\g-ray pulsar. A search for \\g-ray pulsations using EGRET data didn't reveal the pulsar \\citep{Ziegler2008}. The characteristics of this source in X-rays also pointed to it as a pulsar \\citep{Halpern2004}. In fact \\cite{Halpern2004} called the source the ``pulsar'' although no pulsations were detected from this source using FFT searches on $\\sim26$ ks of \\x data \\citep{Slane2004} or in radio with the Green Bank Telescope \\citep{Halpern2004}. Very deep searches for a counterpart for \\r in optical and radio resulted only in upper limits \\citep{Halpern2004}. By correlating X-ray images from \\C with those in the optical waveband of the \\c region, \\citet{Halpern2004} gave the most accurate position of this source to date ((J2000.0) 00$^\\mathrm{h}$07$^\\mathrm{m}$1\\fs56,+73\\degr03\\arcmin08.1\\arcsec) with an accuracy of $\\sim 0\\arcsec.1$. The discovery of the pulsar didn't come until the launch of \\f. During its commissioning stage the \\lsp discovered a 315.87 ms pulsar at the location of RX J0007+7303 \\citep{CTA1}. The discovery of the pulsar in \\g-rays prompted a long, $\\sim130$ ks, \\x observation to search for pulsations from this source (PI: Caraveo 2008, ObsID: 06049401). Using this \\x data set along with the timing model from the \\lsp \\citep{BSP} X-ray pulsations from this source in X-rays were finally detected \\citep{Lin2010,Caraveo2010}.  This paper reports further analysis on the LAT data and related observations of the \\c SNR and its \\g-ray pulsar, extending the initial report of the pulsar discovery \\citep{CTA1}.  The new developments contribute to a characterization both of the pulsar and its relation to the associated extended source, exploiting two years of data now accumulated by \\f.  New developments fall in three distinct areas.  First, timing analysis of the pulsar spanning the two years shows a glitch near the middle of the interval at MJD 54952.652 (1 May 2009), with relative change in pulse frequency $\\Delta\\nu/\\nu \\sim 6 \\times 10^{-7}$.  This permits careful comparison of the pulsar characteristics before and after the glitch.  Second, we present new \\g-ray results concerning the relationship of the pulsar to the surrounding remnant, the first detection of an extended source in the off-pulse emission. Finally, we summarize the multi-wavelength picture of the source. From this perspective the outstanding characteristic of the neutron star in CTA1 is that it appears cool for its inferred age. We reconsider both the temperature and the age estimates, and then relate this to pulsar characteristics established from timing LAT \\g-rays.  In earlier work \\citep{Halpern2004} the pulsar in \\c has been assigned an inferred mass exceeding 1.42 $M_\\sun$ even though it is not in a binary, because the larger mass should accelerate cooling. Even from initial timing solutions it was clear that this pulsar has a very strong magnetic field, though weaker than that of a magnetar, hence it is atypical in two respects. ",
        "conclusions": "\\subsection{Geometrical Constraints from Light Curve Modeling} \\label{sec:LC_modeling} The pulsar emission mechanism is not well understood, and the distributions of pulsar magnetic inclination angles, efficiencies, and other physical characteristics are largely unknown. There are several competing emission models, of which the resulting pulse profiles depend on the geometry of the system, defined by the emission zone location and size, observer viewing angle, and magnetic inclination angle. Fitting the profiles of these models to observed light curves can provide insight on the true emission and viewing geometries, for example leading to better constraints on luminosity and efficiency through calculation of the flux correction factor $f_{\\Omega}$ (equation 4 of \\cite{Watters09}). The geometry of a given system likely also contributes to the detectable presence or absence of radio and perhaps X-ray emission. For example, if the inclination and observer angles are very different, we could see a $\\gamma$-ray-only pulsar, as the narrower radio beam may not cross our line of sight. To determine the geometry of the CTA 1 pulsar and distinguish between emission models, we compared the LAT light curve of PSR J0007+7303 with the predicted light curves from geometrical representations of two standard high-energy pulsar emission models, the outer gap (OG) \\citep{Romani1995} and slot gap (SG, also referred to as the two-pole caustic or TPC) \\citep{Muslimov04} models. These models were considered within the context of the vacuum retarded dipole magnetic field.  The light curves were simulated as in \\cite{Dyks2004}, with the geometry modified to represent each emission region. The OG lies along the last open field lines between the null charge surface and the light cylinder, $R_{\\mathrm{lc}} = c/\\Omega$, with $r_{\\mathrm{cyl}} \\, \\leq \\, 1.0R_{\\mathrm{lc}}$ as large as possible for each geometrical configuration. The SG extends from the surface to $r_{\\mathrm{cyl}} = 0.95R_{\\mathrm{lc}}$. In both geometries, the maximum emission altitude $r_{\\mathrm{max}}$ was treated as a model parameter. We note that our representation of the SG differs from that of the TPC model in \\cite{Romani2010}, as we allow emission out to much higher altitudes (their emission is cut off at $r_{\\mathrm{cyl}} = 0.75R_{\\mathrm{lc}}$). The TPC geometry is therefore treated as a subset of the SG, rather than being considered as a completely separate model. The SG model has emission throughout the gap, while the OG emission occurs only along the field line at the gap's innermost edge. In the simulations, the vacuum retarded dipole field is assumed in the observer's frame and then transformed to the co-rotating frame (CF). Photons are emitted tangent to the field in the CF, prior to the aberration calculation. A constant emissivity is assumed along the field lines in the CF. The special relativistic aberration leads to a bunching in pulse phase, or caustics, on the trailing side of the pulse, producing peaks in the light curvyes. For a given inclination angle $\\alpha$, gap width $w$ in open volume units ($r_{\\mathrm{ovc}}$, as in \\cite{Dyks2004}), and maximum emission radius $r_{\\mathrm{max}}$ in units of $R_{\\mathrm{lc}}$, the code produces light curves at all observer angles $\\zeta$.  The LAT light curve of PSR J0007+7303 has 32 bins and a background count level of 195 or 246 counts/bin, depending on whether the off-peak emission discussed in \\S2.2 is (respectively) magnetospheric or due to a PWN. We compared this light curve with light curves simulated from a representative sample of all possible geometries and rebinned to match the observed light curve. Our models had five free parameters, $\\alpha$, $\\zeta$, $w$, $r_{\\mathrm{max}}$, and a phase shift $\\Delta \\phi$ introduced in order to best match the LAT light curve. We considered values of $\\alpha$ between $0^{\\circ}$ and $90^{\\circ}$ and $\\zeta$ between $0.5^{\\circ}$ and $89.5^{\\circ}$, each with resolution of $1^{\\circ}$; values of width $0 \\leq w \\leq 0.3$ with resolution 0.01$\\,r_{\\mathrm{ovc}}$; and maximum emission radii $0.7 \\leq r_{\\mathrm{max}} \\leq 2$ for the SG and $0.9 \\leq r \\leq 2$ for the OG, with resolution 0.1$\\,R_{\\mathrm{lc}}$ in both cases. For each model, the emission altitude is limited by Min($r_{\\mathrm{max}}, r_{\\mathrm{cyl}}$)---$r_{\\mathrm{max}}$ therefore may be larger than $r_{\\mathrm{cyl}}$, but emission ceases at the cylindrical radius if not at the maximum emission radius. The phase shift $0^{\\circ} \\leq \\Delta \\phi \\leq 360^{\\circ}$ is arbitrary because PSR J0007+7303 is radio-quiet within the flux limits achieved thus far by radio telescopes; were the pulsar radio-loud, the shift would be constrained to be at most equal to the phase lag between the radio and $\\gamma$-ray peaks \\citep{Dyks2004}.   To find the best-fit parameters of each model, we used a Markov Chain Monte Carlo (MCMC) maximum likelihood routine as described in \\cite{Verde2003} that explored the parameter space. We calculated the $\\chi^2$ between the LAT light curve and model light curve and used Wilks' Theorem, $\\Delta\\,$ln$(L) = -\\Delta \\chi^2/2$, to guide the Markov chains toward regions of high likelihood. We also used the $\\Delta \\chi^2$ test to find $3\\sigma$ confidence intervals on each parameter. For this pulsar, the best-fit light curve has a very large $\\chi^2$. This is expected, as a) the pulsar is bright and its light curve has relatively small error bars and b) the light curve simulations result from geometrical models of possible but simplified pulsar magnetospheres, rather than from a well-understood physical model. When finding confidence intervals in steps of $\\Delta \\chi^2$ from a large initial $\\chi^2$ value, the intervals will appear artificially small. We therefore re-scale the $\\chi^2$ values found for all sets of parameters in the MCMC chains so that the minimum reduced $\\chi^2 = N_{\\mathrm{dof}} = 27$ (the scale factor is therefore $N_{\\mathrm{dof}} / \\chi^2_{\\mathrm{min}}$). The confidence intervals are then found with these modified $\\chi^2$ values. In this way, we find parameter ranges that may give rise to the high-energy light curve of PSR J0007+7303.  For the case where the off-peak emission is assumed to be magnetospheric (corresponding to a background level of 195 counts/bin), the best-fit parameters are ($\\alpha$, $\\zeta$, $w$, $r$, $\\Delta \\phi$) = (6$^{\\circ}$, 74.5$^{\\circ}$, 0.04$r_{\\, \\mathrm{ovc}}$, 1.0$\\,R_{\\mathrm{lc}}$, 2$^{\\circ}$) with $\\chi^2/27 = 22.5$ for the OG model and (8$^{\\circ}$, 69.5$^{\\circ}$, 0$\\, r_{\\mathrm{ovc}}$, 1.9$\\,R_{\\mathrm{lc}}$, -4$^{\\circ}$) with $\\chi^2/27 = 11.3$ for the SG model. Almost identical parameters are found for the case where the off-peak emission is assumed to originate from a PWN. The confidence intervals, $f_{\\Omega}$ values, and reduced $\\chi^2$ for the best fit in each interval are given in Table~\\ref{table:model_fits}. The absolute best fit is found with the OG model using the higher background level. Figure~\\ref{fig:model_fits} shows the LAT light curve superposed with the best-fit model light curves, as well as reduced $\\chi^2$ contours in $\\alpha$ and $\\zeta$ with $w$ and $r$ fixed at the best-fit values. To be complete, we also fit the light curve with the TPC model, the results of which are not shown here; we find this lower-altitude subset of the SG model cannot reproduce the sharp double peaks with the correct spacing anywhere in parameter space. We therefore find that both outer magnetosphere models produce light curves consistent with that which is observed.  For magnetospheric off-peak emission, the SG is preferred over the OG; its ability to reproduce the light curve is seen clearly by comparing the red curves in panels ({\\textit a}) and ({\\textit c}) of Figure~\\ref{fig:model_fits}. The reverse is true under the assumption of off-peak emission from a wind nebula, shown by the blue curves. Under the assumption of magnetospheric off-pulse emission, the SG misses some emission in the wings and has too high a background level; this effect is magnified when we assume instead a PWN origin of this emission. Such details may be modified with a more physical emission model, for example by including azimuthal asymmetry in the accelerating electric field from offset polar caps, which leads to a decrease in the off-peak emission \\citep{Harding2011}. The OG characteristically does not reproduce the wings or higher emission in the off-peak phases when the background is derived assuming pulsar emission in the off-peak; by definition it matches the background well when the background is instead found assuming the emission is not from the pulsar itself.   Regardless of the background counts used, the best-fit values of $\\alpha$, and $\\zeta$ for the OG and SG models are far apart, with $|\\zeta-\\alpha| >$60 degrees. The radio beam would need a width of that order in order for the radio pulse to be detectable. For the parameters of \\psrsp a model estimate of beam width is $<$ 10 degrees \\citep{Story07}. Thus the preference for OG and SG model fits over that of TPC, along with the fitted parameter values, self-consistently offers a satisfactory explanation for non-detection of any radio pulse. A deep search using GBT \\citep{Halpern2004} yielded an upper limit that remains the lowest limiting flux for any pulsar position. The upper limit on the pseudo-luminosity of $L_{1400} \\simeq 0.02 $ mJy kpc$^2$ is lower than the lowest pseudo-luminosity measured for a radio pulsar ($L_{1400}\\,\\simeq \\,0.035$ mJy kpc$^2$ for PSR J1907+0602 \\citep{1907}) and is thus restrictive enough to support the radio-quiet designation. Conversely, if one were to adopt the TPC model for this source then the nearly-equal values of $\\alpha$ and $\\zeta$ would suggest that a radio pulse should have been found.  X-ray pulsations were recently detected from the CTA 1 pulsar using the same LAT ephemeris used in the timing and light curve analysis of this paper \\citep{Lin2010,Caraveo2010}.  We did not include any X-ray information in our fits, but can consider whether or not our results make sense in the multiwavelength picture. The peak in the X-ray light curve shows a significant thermal component, suggestive of a hot spot on the surface. The authors estimate the size of the hot spot to be $\\sim 100\\,$m in radius, larger than the polar cap in the case of a dipole model for the pulsar; it is not clear from the data where the hot spot is located. The X-ray peak occurs at $\\phi \\sim 0.25$--0.3 ($\\sim 90^{\\circ}$--110$^{\\circ}$) prior to the first peak in the $\\gamma$-ray light curve. Our HE models predict a similar location of the magnetic pole: the first $\\gamma$-ray peak is at $90^{\\circ}$, and in all our model fits the pole lies $36^{\\circ}$--$94^{\\circ}$ before this peak; taking only the SG and best OG fits places the pole $94^{\\circ}$--$88^{\\circ}$. Our best-fit models are therefore consistent with the thermal X-ray emission originating near the magnetic pole.  \\subsection{Thermal Component in X-rays and Neutron Star Cooling}  \\psrsp is among the neutron stars (NS) where the theory of cooling confronts observation. It held that distinction before launch of \\f, based on the candidate NS identified at that time \\citep{Halpern2004}, and it continues to be treated as an exceptional case in the cooling literature \\citep{Page2009}.  Those references cite additional literature on heat loss models which describe how observable surface temperature or thermal flux declines with age.  X-ray or extreme ultra violet surface emission in young, hot NS provides the observational test.  The NS age and surface temperature must be known or constrained, that is, a limit may be useful if corrections can be characterized at least as to sign if not magnitude.  Only comparatively young NS within a few kpc provide useful constraints. The only pulsar both younger and closer than \\psrsp is the Vela pulsar.  Before \\f, the age of \\psrsp was equated to the estimated CTA1 SNR age and thermal flux was estimated from the X-ray emission of the identified point-like NS candidate \\citep{Halpern2004}. From this it was understood that the putative pulsar in CTA1 was cool for its age, in comparison with both theory and other young NS.  This led to theoretical effort to understand why this particular NS had cooled rapidly.  \\fermis results impact context information that enters into the analysis of the cooling question.  The pulsar ephemeris affects the X-ray analysis used to extract the thermal flux.  The \\fermis results have not greatly altered best-estimates of key quantities but provide a new network of interlocking constraints that give greater robustness of the observational context information, as follows.  (1) The spin period $P$ and its derivative $\\dot{P}$ give spindown age of $\\sim$13,900 yr for the pulsar.  This age estimate is well within the broader range of historical estimates for the SNR age, 5,000-15,000 yr, and nearly identical to one pre-launch estimate of $1.3 \\times 10^4 \\, d_{1.4} $ yr \\citep{Slane2004}.  The formal propagated error in the spindown age estimate is negligible.  The principal issue is validity of the dipole approximation formula but it clearly reinforces the age estimates derived other ways. (2) The same observations ($P$, $\\dot{P}$) also yield spindown energy loss independently of distance. There are no inconsistencies between this energy budget and that of the pulsar plus SNR, using the accepted distance, D, of $1.4 \\pm 0.3$ kpc, hence the pre-\\fermis distance estimate used is not in conflict with new information and the distance scale factor in the age estimate just cited is consistent with unity. (3) The pulsar ephemeris has now been used to detect X-ray pulsations using XMM \\citep{Lin2010, Caraveo2010}.  The X-ray spectrum can be divided into pulsed and unpulsed components. While total X-ray flux is consistent with levels measured in pre-\\f\\ work, the part now assigned to a DC thermal component emitted from the surface is reduced by at least a factor that equates to the unpulsed fraction.  Since D is unchanged, the revised bound on NS surface temperature remains as in earlier estimates or perhaps even slightly reduced.  One may consider modifications due to the interstellar absorbing column and to the neutron star atmosphere causing the surface flux to depart from blackbody.  These considerations properly enter into many other NS used for comparison with cooling and it is beyond scope of this paper to reanalyze the entire X-ray discussion. It is left to future research but provisionally the pre-launch thermal luminosity stands. (4) In this same connection a further point to note is that the off--pulse \\g-ray flux, discovered and quantified in \\S \\ref{sec:offpulse} of the present paper, cannot be thermal flux from the neutron star and must be taken as magnetospheric or PWN emission. In either case, that same magnetosphere or PWN component could extend downward in energy to X-rays, and in principle could provide further downward adjustment to the thermal X-ray flux from the star.  This is left as an adjustment of unknown magnitude but known sign; it can only require the star to cool faster than the estimate obtained by neglecting it. (5) From $P$ and $\\dot{P}$ one also obtains an estimate of the stellar dipole field, B $\\sim10^{13}$G. \\psrsp therefore falls among the most highly magnetic neutron stars, magnetars excepted. Cooling scenarios whereby a strong magnetic field could modify the cooling curve have been described in earlier literature (for a summary see for example \\cite{Yakovlev2001}).  The high value of B now established for this pulsar means mechanisms whereby high B enhances cooling may merit further attention.  (6) The previous section shows how model fits for OG and SG models favor small $\\alpha$ and large $\\zeta$, and could explain the absence of radio pulsations.  The first four items in the list mean that if \\psrsp was an outlier relative to models before \\textit{Fermi} and relative to pre-\\textit{Fermi} theoretical understanding, it has become slightly more egregious relative to that prior theoretical understanding. However theory has also advanced, largely from observation of cooling in the central star in Cas A \\citep{Page2009}.  Data from that source are now the most definitive and constraining of any un-recycled NS. Fitting Cas A has given prominence to NS cooling models involving Cooper pairing contributions.  Points (5) and (6) in the list regarding magnetic field strength and geometry may provide guidance to theory.  Sufficiently strong magnetic fields can affect cooling \\citep{Yakovlev2001} and might affect heat flow in the star.  Also the particular geometry in \\psrsp with low $\\alpha$ and high $\\zeta$ that operates against radio pulse detection could also lead to a misleadingly low thermal X-ray flux if the strong field conveyed internal heat flow preferentially to the magnetic polar regions. Then, a Lambertian emission pattern from the hotter poles would be anisotropically beamed away from an observer at high $\\zeta$.  Further X-ray observations, combined with multi-wavelength analysis applied and field geometry modeling may shed further light on the thermal luminosity.  This could be undertaken comparatively with other young, nearby pulsars such as the \u201cDragonfly\u201d pulsar, PSR J2021+3651, which has estimated age 17 kyr, distance 2.1 kpc, dipole field $3 \\times 10^{12}$ G, and has radio pulses.  Such comparisons might bring out the role of the magnetic field strength and geometry in NS cooling."
    },
    "1107/1107.3895_arXiv.txt": {
        "abstract": "F00183-7111 is one of the most extreme Ultra-Luminous Infrared Galaxies known. Here we present a VLBI image which shows that F00183-7111 is powered by a  combination of a radio-loud Active Galactic Nucleus surrounded by vigorous starburst activity. Although already radio-loud, the quasar jets are only 1.7 kpc long, boring through the dense gas and starburst activity that confine them. We appear to be witnessing this remarkable source in the brief transition period between merging starburst and radio-loud ``quasar-mode'' accretion. ",
        "introduction": "\\label{intro} Ultraluminous infrared galaxies (ULIRGs) are a class of galaxy with a bolometric luminosity in excess of $10^{12}$ $L_{\\odot}$ which were first uncovered by the IRAS satellite more than 20 years ago \\citep{Aaronson:1984p29003, Houck:1985p29887, Allen:1985p29041}. A variety of observations suggest that they represent a transitional stage in which gas-rich spirals are merging to form a dusty quasar \\citep{Armus:1987p29090,Sanders:1988p29218, Veilleux:2002p29465, Spoon:2009p29237}.   According to the scenario by \\citet{Sanders:1988p29218}, the merger of two normal spirals fuels a pre-existing quiescent black hole and triggers a powerful nuclear starburst.This intense starburst activity causes the high infrared luminosity, and generates strong starburst-driven winds which will eventually blow away the enshrouding dust and lay bare the quasar core, depleting  the dust and gas to form an elliptical galaxy \\citep{Dasyra:2006p29091, Dasyra:2006p29094}.   This scenario is supported by merger simulations \\citep[e.g.][]{DiMatteo:2005p29098, Hopkins2005}, which show that  the supermassive black hole grows by accretion while enclosed by obscuring dust, and then sheds its obscuring cocoon, deposited by the merger, through outflows driven by powerful quasar winds \\citep{Balsara:1993p28980}. This activity ceases when the fuel supply to the central regions is exhausted, starving both the active galactic nucleus (AGN) and the star forming activity. This scenario is also supported by the observational result that AGN activity increases with increasing bolometric luminosity \\citep{Genzel:1998p29109, Tran:2001p29490, Lonsdale2006}.   In the low-redshift Universe, star formation is dominated by M82-type starburst galaxies, and less than 50 ULIRGs are known at $z<0.1$. However, at higher redshifts  the cosmic star formation rate is dominated by increasingly vigorous starburst galaxies. At $z=0.5$ ULIRGs are far more numerous than in the local Universe, and dominate the cosmic star formation rate \\citep{Blain:2002p28984, LeFloch2005, Lonsdale2006, Wilman:2008}. They appear to represent a significant step in the evolution of galaxies, so that understanding their nature is crucial to understanding galaxy evolution.  However, the many magnitudes of extinction to their nuclei make it difficult to determine whether their dominant power source is accretion onto a supermassive central black hole, or a colossal starburst. The current consensus \\citep[e.g.][]{Risaliti10, Veilleux09} is that most ULIRGs are primarily powered by a starburst, although an AGN is dominant in a significant fraction. However, this view is based largely on optical and near-infrared observations, which penetrate only the outer layer, and tell us little of the nucleus.   IRAS F00183-7111 (also known as IRAS 00182-7112, and which we shall hereafter refer to as 00183) is the most luminous ULIRG discovered in the IRAS survey, and one of the most luminous ULIRGs known. At a redshift of 0.3276 \\citep{Roy:1997p29162}, it has a bolometric luminosity of $9 \\times10^{12}$ \\Lo \\citep{Spoon:2009p29237}, most of which is radiated at far-infrared wavelengths, and a linear scale of 1 mas = 4.7 pc ($H_{0} = 71$ km s$^{-1}$ Mpc$^{-1}$). It is the most luminous of the extreme ULIRGs studied by \\citet{Armus:1987p29090}, and also has the largest FIR excess,  with $L_{\\rm{FIR}} / L_{\\rm{B}} = 360$.  For comparison, Arp 220 has $L_{\\rm{FIR}} / L_{\\rm{B}} = 150$, and spectacular starburst systems like M82 and NGC 3256 have $L_{\\rm{FIR}} / L_{\\rm{B}} = 10$, so this is a heavily obscured object. Near-infrared imaging by \\citet{Rigopoulou:1999p29130} (Figure \\ref{fig:figIR}) appears to show a disturbed morphology and a single nucleus, although the many magnitudes of extinction, even at near-IR wavelengths, mean that this image may be showing an outer shell of the galaxy, rather than the nucleus. CO emission has also been detected from 00183 with an inferred molecular gas mass $> 2.4 \\times 10^{10} $ \\Msun \\citep{Norris11}, confirming that this object is a vigorous star-forming galaxy.  Optically \\citep{Drake:2004p30337}, it appears as an unremarkable smudge 20 arcsec across \\citep{Heckman1990}, but long-slit spectroscopy  shows bright ionised gas extending 50 kpc east and 10 kpc west of the nucleus, with highly disturbed kinematics.  Line widths are $> 600$ km s$^{-1}$ over a 30 kpc-diameter region, reaching 1000 km s$^{-1}$ in the region 3 - 7 kpc east of the nucleus.  Line profiles are double-peaked with a splitting of 300 - 400 km s$^{-1}$ in the region 10 kpc east, and the profiles are strongly blue-shifted with respect to the nucleus out to 25 kpc E.  The line ratios and widths are typical of Sy2 galaxies, and the kinematics suggest a powerful outflow.  \\cite{Spoon:2009p29237}  observed unusually wide [Ne II] and [Ne III] profiles in 00183 -- the widest [Ne II] and second widest [Ne III] lines in a sample of 82 ULIRGs observed with Spitzer-IRS \\citep{Spoon09}. The lines may trace strongly disturbed gas, resulting from interaction of the radio jets with the ISM. However, the strong dust extinction means that these are characteristics of the outer parts of the galaxy, rather than the nucleus. The size, velocity, and morphology of the outflow agrees well with the super-wind model of \\citet{Heckman1990}, in which the optical emission comes from a thin shell of shock-heated ambient gas swept up by a wind driven by the overpressure produced by the starburst.  \\citet{Spoon:2004p28328, Spoon2007} fail to detect several common mid-infrared AGN tracers in their Spitzer-IRS high resolution spectrum. Furthermore, \\citet{Spoon:2004p28328} detect a deep 9.7 $\\mu$m silicate absorption feature, and  strong  4.7 $\\mu$m  CO absorption. All of these are unusual for classical Seyfert galaxies and quasars \\citep{Spoon2007, Hao:2007p29931} suggesting that the line of sight to the nucleus is heavily obscured.  However, radio observations are unaffected by the heavy dust obscuration that blocks our view at shorter wavelengths, and allow us to see into the nucleus to look for typical AGN-related core-jet structures with high brightness temperature, to distinguish from diffuse, lower $T_{\\rm{b}}$ starburst emission. At the centre of 00183 lies a strong radio source \\citep{Roy:1997p29162}, whose radio luminosity is ten times higher than would be expected from a star-forming ULIRG of 10$^{13}$ \\Lo. Instead, this radio luminosity is attributable to a radio-loud AGN buried at the core of 00183, with a radio luminosity of $L_{4.8GHz}=3\\times10^{25}$ W Hz$^{-1}$ \\citep{Roy:1997p29162} which places it within the regime of high luminosity (FRII-class) radio galaxies. This AGN is invisible at optical and near-infrared wavelengths because of the dense dust galaxy surrounding it, evidence for which includes  the deep 9.7 $\\mu$m silicate absorption feature \\citep{Tran:2001p29490, Spoon:2004p28328}. However, the AGN is confirmed by the detection of a 6.7 keV FeK line (Fe XXV) with a large equivalent width, indicative of reflected light from a Compton thick AGN \\citep{Nandra:2007p29707}.  In a sensitive 12-hr ATCA synthesis observation of 00183 at 4.8 GHz, Roy et al. \\citep[unpublished data cited by ][]{Roy:1997p29162} found an unresolved core of 108 mJy and no lobes brighter than 30 K ($3 \\sigma$).  The core has an extraordinary radio luminosity of $L = 3\\times10^{25}$ W Hz$^{-1}$, typical of powerful radio galaxies, which makes the object $\\approx 40$ times more radio-loud than expected for the normal radio-FIR correlation.  To produce this radio luminosity using supernovae, as in Arp 220, would require an implausible  20000 luminous radio supernovae (i.e. similar to RSN 1986J) radiating simultaneously.    However, the ultimate distinction between radio-loud AGN and star formation is best achieved using Very Long Baseline Interferometry (VLBI). While clumps of radio supernovae can emulate low-luminosity AGN, high luminosity VLBI detections are nearly always produced by AGN \\citep{Kewley:2000p30027}. Here we present the results of a VLBI experiment to probe the radio core of 00183. ",
        "conclusions": "Our VLBI image of the extreme ULIRG F00183-7111 shows  a classical core-jet radio-loud ($ > 10^{25}$ W Hz$^{-1}$ ) AGN, which has the luminosity and morphology of an FRII galaxy, except that the total length of the source is only 1.7 kpc. This core is surrounded by a vigorous starburst accompanied by some $ 2.4 \\times 10^{10} $ \\Msun\\ of molecular gas. The many magnitudes of dust extinction  accompanying this starburst totally obscure any sign of this AGN at optical or infrared wavelengths.  We have placed this source in an evolutionary sequence defined by the standard quasar-mode/radio-mode model, and find that it represents the initial stage of the formation of a quasar. It is in the earliest throes of quasar-mode accretion, with the powerful radio jets still confined by the dense cool gas from the galaxies that merged to form this system.  In this respect it is very similar to the classical CSS/GPS sources, differing only in that its host galaxy appears to be at an earlier stage of post-merger evolution than most CSS/GPS hosts. Eventually the radio jets will drill their way out of the gas, at the same time heating and disrupting the gas. As a result, the current vigorous burst of star formation will be extinguished, marking the transition of this source from a star-forming ULIRG with a ``quasar-mode''  AGN at its centre to a powerful radio-loud quasar."
    },
    "1107/1107.1542_arXiv.txt": {
        "abstract": "We present an observational study towards the young high-mass star forming region G23.44-0.18 using the Submillimeter Array. Two massive, radio-quiet dusty cores MM1 and MM2 are observed in 1.3 mm continuum emission and dense molecular gas tracers including thermal CH$_3$OH, CH$_3$CN, HNCO, SO, and OCS lines. The $^{12}$CO (2--1) line reveals a strong bipolar outflow originated from MM2. The outflow consists of a low-velocity component with wide-angle quasi-parabolic shape and a more compact and collimated high-velocity component. The overall geometry resembles the outflow system observed in the low-mass protostar which has a jet-driven fast flow and entrained gas shell. The outflow has a dynamical age of $6\\times10^3$ years and a mass ejection rate $\\sim10^{-3}~M_{\\odot}$ year$^{-1}$. A prominent shock emission in the outflow is observed in SO and OCS, and also detected in CH$_3$OH and HNCO. We investigated the chemistry of MM1, MM2 and the shocked region. The dense core MM2 have molecular abundances of 3 to 4 times higher than those in MM1. The abundance excess, we suggest, can be a net effect of the stellar evolution and embedded shocks in MM2 that calls for further inspection. ",
        "introduction": "The outflows take place in high-mass star forming regions at very early evolutionary stages~\\citep[e.g.][]{birk06,beu07,long11}. They inject large amount of hot gas and kinetic energy into the molecular cloud, causing intense shock waves that dramatically alter the chemistry of the surrounding environment~\\citep{dish98}. However, the morphological and dynamical properties of the outflows on smaller scales, as well as their chemical effect to the young high-mass stellar cores are still to be further characterized.  The high-mass star forming region G23.44-0.18 (G23.44 hereafter) was previously detected as a group of strong CH$_3$OH masers~\\citep{walsh98} which indicate the presence of young massive stars. It has a trigonometric-parallax distance of 5.88 kpc~\\citep{brun09}. In this letter, we report an observation of the massive dusty cores MM1 and MM2 associated with the CH$_3$OH masers, as well as a high-velocity, intense bipolar outflow revealed in CO (2--1). The outflow is causing a shock emission and may be affecting the chemistry in MM2. Section 3 presents the observational results. A further discussion on the outflow is given in Section 4.1 and 4.2. The properties of the dusty cores are discussed in Section 4.3. A summary is given in Section 5. ",
        "conclusions": "\\subsection{The physical properties of the outflow} The presence of intense bipolar outflow in G23.44 suggests that the central stars are likely being formed via a disk-mediated accretion~\\citep{zhang07}. The HVC is well collimated along the outflow axis, while the LVC, especially its red wing, has a more opened parabolic shape with a decreasing brightness from the edge to the center. The overall geometry of the outflow is reminiscent of the outflow system in the low-mass protostar HH211 \\citep{gueth99}. HH211 consists of a high-velocity jet-driven flow and a entrained low-velocity gas shell. The G23.44 outflow may represent a scaled-up version of this system, despite its collimation and shell structure being much less perfect than HH211. However, the size of the outflow in HH211 is only $\\sim10^4$ AU, while the G23.44 outflow ($10^5$ AU) has a 10-time larger spatial extent. On such a large scale, the outflow structure might begin to deteriorate due to increased turbulence in its environment.  \\begin{table} {\\small \\centering \\begin{minipage}{80mm} \\caption{Abundance of the detected molecules.} \\begin{tabular}{lccc} \\hline\\hline species                        &  $N_{\\rm x}$        & $f_{\\rm x}$           &  $f_{x,2/1}^a$  \\\\ (X)                            &  (cm$^{-2}$)        & \\quad                 &  \\quad \\\\ \\hline \\multicolumn{4}{l}{MM1 $(-1'',10'')$} \\\\ C$^{18}$O                    &  $3.16\\pm0.05(16)$    & $1.7\\pm0.1(-7)$         &   $-$ \\\\ CH$_3$OH                     &  $2.5\\pm1.2(15)$      & $1.3\\pm0.6(-8)$         &   $-$ \\\\ CH$_3$CN                     &  $5.2\\pm0.6(14)$      & $2.7\\pm0.5(-9)$         &   $-$ \\\\ HNCO                         &  $3.8\\pm0.2(14)$      & $2.0\\pm0.4(-9)$         &   $-$ \\\\ SO                           &  $2.5\\pm0.2(14)$      & $1.3\\pm0.3(-9)$         &   $-$ \\\\ OCS                          &  $3.7\\pm0.2(14)$      & $1.9\\pm0.3(-9)$         &   $-$ \\\\ \\hline \\multicolumn{4}{l}{MM2 $(1'',-3'' )$} \\\\ C$^{18}$O                    &  $7.30\\pm0.07(16)$     & $2.9\\pm0.1(-7)$        &   $1.7\\pm0.2$  \\\\ CH$_3$OH                     &  $9.8\\pm0.8(15)$       & $3.9\\pm0.9(-8)$        &   $3.0\\pm1.5$  \\\\ CH$_3$CN                     &  $1.6\\pm0.2(15)$       & $6.3\\pm0.7(-9)$        &   $2.3\\pm0.8$  \\\\ HNCO                         &  $2.20\\pm0.05(15)$     & $8.8\\pm1.5(-9)$        &   $4.4\\pm1.4$  \\\\ SO                           &  $1.45\\pm0.06(15)$     & $5.8\\pm0.9(-9)$        &   $4.5\\pm1.4$  \\\\ OCS                          &  $1.20\\pm0.05(15)$     & $4.8\\pm0.8(-9)$        &   $2.5\\pm0.7$  \\\\ \\hline \\multicolumn{4}{l}{outflow-shock $(-11'',10'')$} \\\\ C$^{18}$O                   &  $2.4\\pm0.07(16)$       & $=2.0(-7)^b$           &  $-$    \\\\ CH$_3$OH                    &  $2.8\\pm0.5(15)$        & $2.3\\pm0.5(-8)$        &  $-$    \\\\ HNCO                        &  $4.1\\pm0.2(14)$        & $8.3\\pm0.4(-10)$       &  $-$    \\\\ SO                          &  $3.5\\pm0.3(14)$        & $3.1\\pm0.3(-9)$        &  $-$    \\\\ OCS                         &  $8.0\\pm0.3(14)$        & $6.6\\pm0.3(-9)$        &  $-$    \\\\ \\hline \\end{tabular}  $a.${ The abundance ratio of MM2 to MM1.} \\\\ $b.${ Assumed as the typical value in ISM. Abundances of the other molecules are then calculated by comparing to $N({\\rm C^{18}O})$.} \\end{minipage} } \\end{table}  The CO outflow, especially the LVC also exhibits a weak emission feature in opposite direction to the major one, i.e. a blueshifted emission in the NW and a redshifted one in the SE. This can be a second flow by chance aligned in the same position angle with the major one. Alternatively, once the outflow direction is close to the plane of the sky and its inclination angle becomes smaller than the opening angle $(\\sim30^{\\circ})$, its front and back sides will exhibit opposite Doppler shift. This scenario is well demonstrated by~\\citet[][Figure 10 therein]{yen10}. We speculate the second case to be more plausible.  Assuming a local thermal equilibrium (LTE) and low optical depth for the $^{12}$CO line-wing emission, the CO column density and the outflow mass can be inferred following the approach of~\\citet{gard91}. The excitation temperature of the outflow is speculated to be slightly lower than in the dense core MM2 ($\\sim70$ K, Section 4.3). A value of $T_{\\rm ex}=50$ K is adopted here. The outflow mass is estimated in each 1 km s$^{-1}$ interval, then the total mass, momentum, and kinetic energy are calculated from $M=\\Sigma_v m(v)$, $P=\\Sigma_v m(v) v$, and $E_k=\\Sigma_v m(v) v^2/2$. In calculation, we assume an inclination angle of $\\theta=20^\\circ$ to correct the velocity projection. The outflow turns out to have $M\\simeq5.5~M_{\\odot}$, $P\\simeq360~M_{\\odot}$ km s$^{-1}$, $E_k\\simeq9\\times10^{46}$ erg.  Each of the blue and red wings in the HVC is resolved into two major clumps (Figure 2b, lower panel) denoted B1, B2 and R1, R2. These clumps in the outflow may indicate an episodic mass ejection~\\citep[e.g.][]{qiu09}. The average central distance is $24''$ (0.68 pc) for the pair of B1-R1, and $5.0''$ (0.14 pc) for B2-R2. The two pairs both show a reasonable geometrical symmetry. However, the clumps are not exactly aligned across the MM2 center, instead with an offset to its west. This may indicate that the driven agent, or the stellar disk system is also displaced from the MM2 center. The dynamical age of the outflow $t_{\\rm dyn}$ is estimated by dividing its spatial extent with the typical velocity ($v=40/\\sin\\theta$ km s$^{-1}$). That yields $t_{\\rm dyn}=6.1\\times10^3$ years for the B1-R1 pair and $t_{\\rm dyn}=1.3\\times10^3$ years for B2-R2. Adopting the first value, we calculate the mass ejection rate to be $\\dot{M}=M/t_{\\rm dyn}\\simeq1.0\\times10^{-3} M_{\\odot}$ yr$^{-1}$. One can see that the outflow is young and performing an intense mass ejection.  \\subsection{Chemical enhancement in the outflow and shock} In both low- and high-mass young stellar objects the outflows have led to a rich chemistry in CH$_3$OH, HNCO, and sulphides~\\citep[][etc.]{blake87,tak03,jog04,rod10}. In G23.44, the SO and OCS emission shows a noticeable emission clump at offset=$(-11'',10'')$ (Figure 2a).  The clump is almost perfectly aligned with the HVC axis of the CO (2--1), and reasonably coincides with the peak R2, despite being closer to MM2 center. This strongly suggests that the clump traces a shock within the jet/outflow. The clump is also weakly detected in HNCO and CH$_3$OH (spectra shown in Figure S1). Their column densities are estimated using Equation (1) and (2) in Section 4.3. In calculation a temperature of $T_{\\rm ex}=50$ K [same with that of CO (2--1)] is adopted, and the C$^{18}$O abundance is taken to be a constant of $2\\times10^{-7}$ as the typical ISM abundance which is close to $f_{\\rm C^{18}O}$ measured in MM1 and MM2. The abundances of other molecules are then inferred from their intensity ratio relative to C$^{18}$O (2--1). The abundances of the shocked clump are listed in Table 2.  In the low-mass protostars, the outflows often lead to a rich chemistry. In a few cases, the column density of HNCO and S-bearing molecules can been enhanced for orders of magnitude \\citep[e.g. L1157 and L1448,][]{rod10,taf10}. In comparison, the shocked region at $(-11'',10'')$ shows moderate abundances which are comparable or even slightly lower than in MM2. The enhancement level can be limited by the temperature and density in the shocked region, as well as the efficiency of the previous grain-surface chemistry. In addition, since the outflow has a short $t_{\\rm dyn}$, the shock chemistry may still be undergoing and have not yet attained its maximum level.  The blue lobe of CO $(2-1)$ outflow are devoid of the SO and OCS emissions. As a possible explanation, the S-bearing species (e.g. H$_2$S and CS as their progenitors) would be rather tightly bound to the dust grains~\\citep{blake87}, and and a threshold shock velocity may be required for their desorption. Higher resolution and primary shock products like SiO and H$_2$S may better reveal the shock chemistry.  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.4\\textwidth]{f3.pdf} \\\\ \\caption{\\small (a) The solid line is the CH$_3$CN ($12_K-11_K$) towards the center of MM1 (top panel) and MM2 (bottom panel). The dashed line is the spectrum of the best fit from the LTE gas model. (b) The rotation diagram of the CH$_3$OH lines. The filled circles are for the observed transitions. The vertical bars indicate the $3\\sigma$ errors. The linear least-squares fit is shown as the solid line.} \\end{figure}  \\subsection{The dense core properties} The dust cores MM1 and MM2 coincide with dark patches on the extended 8 $\\micron$ emission (Figure 1). The two cores are slightly connected with each other, suggesting that they might be the fragmentation products from the same natal cloud.  The CH$_3$CN ($12_K-11_K$) lines provide a temperature diagnosis for the cores (Figure 3a). We take an LTE radiative transfer model with a uniform rotation temperature and gas density~\\citep{chen06,zhang07} to reproduce the observed CH$_3$CN spectra at the dust-core center. We found the best fit at $T_{\\rm rot}=65$ K for MM1 and $T_{\\rm rot}=110$ K for MM2. The modeled spectra of $K=0$ and 1 towards MM2 is much lower than the observed line profile while the $K>1$ lines are all well fitted. The excess emission in $K=0$ and 1 lines may largely arise from the cold outer-envelope gas.  The rotation temperature is also evaluated from the thermally excited CH$_3$OH lines at the core center using the rotation diagram method~\\citep{gold99,liu01,biss07}. The column density of the upper level $N_{\\rm u}$ is related to the line intensity in the form \\begin{equation} \\qquad \\qquad \\frac{N_{\\rm u}}{g_{\\rm u}}=\\frac{3k}{8\\pi^3\\nu S\\mu^2}\\int T_{\\rm b} \\, {\\rm d}v \\end{equation} where $g_{\\rm u}$ is the upper-level degeneracy. $S$ is the line strength, $\\mu$ is the dipole moment. $N_{\\rm u}$ depends on the total column density $N_{\\rm T}$ in the form \\begin{equation} \\qquad \\qquad \\ln(\\frac{N_{\\rm u}}{g_{\\rm u}})=\\ln(\\frac{N_{\\rm T}}{Q_{\\rm rot}})-\\frac{E_{\\rm u}}{T_{\\rm rot}} \\end{equation} where the partition function $Q_{\\rm rot}$ can be adopted from~\\citet{blake87}. $T_{\\rm rot}$ is then estimated from a least-square fit (Figure 3b). The derived values are listed in Table 1.  The $T_{\\rm rot}$ derived from CH$_3$OH and CH$_3$CN are remarkably similar for MM1. For MM2, however, the $T_{\\rm rot}$ of CH$_3$OH is lower than that of CH$_3$CN. As a probable reason, the CH$_3$OH lines generally have lower $E_{\\rm u}$ than CH$_3$CN ($12_K-11_K$), then the cooler outer-part gas would contribute more significantly to the CH$_3$OH lines than to the CH$_3$CN (in particular $K=2$ to 6). In addition the CH$_3$OH can be easily released from the dust grains by the shocks while the CH$_3$CN production requires subsequent gas-phase reactions and stellar heating~\\citep{dish98}. CH$_3$CN may therefore be distributed closer to the MM2 center (as seen in Figure 2a), and showing a higher $T_{\\rm rot}$.  For the other species, an independent temperature estimation is currently unavailable. We adopt $T_{\\rm ex}=60$ K for MM1 and $T_{\\rm ex}=70$ K for MM2 assuming an LTE excitation, the column density of the molecules are calculated from their spectral lines towards the dust-core center, using equation (1) and (2).  Assuming optically thin 1.3 mm dust continuum emission, the total dust-and-gas mass can be estimated using $S_{\\nu}=M_{\\rm core}\\kappa_{\\nu}B_{\\nu}(T_{\\rm d})/g D^2$. In the formula $S_{\\nu}$ is the total flux density at 1.3 mm. $g=100$ is the dust-gas mass ratio. $B_{\\nu}(T_{\\rm d})=(2h\\nu^3/c^2)/[\\exp(T_0/T_{\\rm d})-1]$ is the Planck function ($T_0=h\\nu/k$). The dust opacity is set to be $\\kappa_{\\nu}=\\kappa_{\\rm 230 GHz}=0.9$ cm$^2$ g$^{-1}$ of MRN dust grains with thin ice mantles at $n_{\\rm H_2}=10^6$ cm$^{-3}$~\\citep{ossen94}. In the expression of $S_{\\nu}$, using the beam-averaged flux density towards the core center, and dividing the obtained mass with $m_{\\rm H_2}$, one can also calculate the peak column density N(H$_2$) at MM1 and MM2. The derived values are presented in Table 1. MM2 has a comparable but slightly higher N(H$_2$) than MM1.  The molecular abundances are estimated using $f_{\\rm x}=N(X)/N({\\rm H_2})$. The $f_{\\rm x}$ and its ratio between MM2 and MM1 are listed in Table 2. For the high-density tracers, the abundances in MM2 are generally 3 to 4 times higher than those in MM1. The SO shows the highest abundance ratio of 4.5. We note that the temperature adopted for MM2 (70 K) is conservative. A higher $T_{\\rm ex}$ would yield a lower N(H$_2$) and higher $N(X)$ for most species, hence leading to even higher abundances. For instance, adopting the value of $T_{\\rm rot}({\\rm CH_3CN})=110$ K as the $T_{\\rm ex}$ for MM2, we will have $N({\\rm HNCO})=6.7\\times10^{15}$ cm$^{2}$ that is 3 times larger than the current value.  The molecular distribution appears to be influenced by the outflow. Besides the coherent elongation of SO and OCS with the HVC of $^{12}$CO emission (Figure 2b), the CH$_3$OH and HNCO emissions are also largely distributed to the west of MM2. In addition, all the molecular lines have broader line widths in MM2 than in MM1, e.g. The SO line has $\\Delta v_{\\rm FWHM}=11.3$ in MM2 and 3.4 km s$^{-1}$ in MM1. The line widths of MM2 can be broadened by a kinetic energy input from the outflow. Considering its influence to the molecular line profiles, the outflow may also contribute to the higher abundances in MM2. Specifically, there would be underlying shocks which are induced as the outflow is interacting with the envelope and/or infalling gas~\\citep{arce04,wu09}.  At the near- to mid-IR and sub-millimeter wavebands, the two cores have similar flux densities (Table 1). At IRAC bands MM2 is brighter than MM1 at 5.8 and 8 $\\micron$ but weaker at 3.6 and 4.5 $\\micron$, while at 2MASS $J$, $H$, $K_s$ bands, only MM1 is significantly detected (Figure S2). Having a spectral energy distribution (SED) obviously skewed to the shorter wavelengths ($\\lambda<4.5~\\micron$), MM1 may have a more evolved stage for its stellar disk system~\\citep{rob06}. However, since the outflow in MM2 is close to the plane of the sky, it is also possible that a disk/toroid structure lies edge-on in MM2 thus would largely absorb the short-wave emissions. A higher resolution will be helpful in revealing the morphology of the stellar system and underlying shocks in MM2."
    },
    "1107/1107.4367_arXiv.txt": {
        "abstract": "We present new results on the Ly$\\alpha$ emission-line kinematics of 18 $z\\sim2-3$ star-forming galaxies with multiple-peaked Ly$\\alpha$ profiles.  With our large spectroscopic database of UV-selected star-forming galaxies at these redshifts, we have determined that $\\sim$30$\\%$ of such objects with detectable Ly$\\alpha$ emission display multiple-peaked emission profiles. These profiles provide additional constraints on the escape of Ly$\\alpha$ photons due to the rich velocity structure in the emergent line.  Despite recent advances in modeling the escape of Ly$\\alpha$ from star-forming galaxies at high redshifts, comparisons between models and data are often missing crucial observational information.  Using Keck II NIRSPEC spectra of H$\\alpha$ ($z\\sim2$) and [OIII]$\\lambda$5007 ($z\\sim3$), we have measured accurate systemic redshifts, rest-frame optical nebular velocity dispersions and emission-line fluxes for the objects in the sample.  In addition, rest-frame UV luminosities and colors provide estimates of star-formation rates (SFRs) and the degree of dust extinction.  In concert with the profile sub-structure, these measurements provide critical constraints on the geometry and kinematics of interstellar gas in high-redshift galaxies.  Accurate systemic redshifts allow us to translate the multiple-peaked Ly$\\alpha$ profiles into velocity space, revealing that the majority (11/18) display double-peaked emission straddling the velocity-field zeropoint with stronger red-side emission.  Interstellar absorption-line kinematics suggest the presence of large-scale outflows for the majority of objects in our sample, with an average measured interstellar absorption velocity offset of $\\langle \\Delta v_{abs} \\rangle =-230$ km s$^{-1}$.  A comparison of the interstellar absorption kinematics for objects with multiple- and single-peaked Ly$\\alpha$ profiles indicate that the multiple-peaked objects are characterized by significantly narrower absorption line widths. We compare our data with the predictions of simple models for outflowing and infalling gas distributions around high-redshift galaxies. While popular ``shell\" models provide a qualitative match with many of the observations of Ly$\\alpha$ emission, we find that in detail there are important discrepancies between the models and data, as well as problems with applying the framework of an expanding thin shell of gas to explain high-redshift galaxy spectra. Our data highlight these inconsistencies, as well as illuminating critical elements for success in future models of outflow and infall in high-redshift galaxies. ",
        "introduction": "The process described as ``feedback\" is considered a crucial component in models of galaxy formation. Feedback commonly refers to large-scale outflows of mass, metals, energy, and momentum from galaxies, which therefore regulate the amount of gas available to form stars, as well as the thermodynamics and chemical enrichment of the surrounding intergalactic medium (IGM). The evidence for feedback in high-redshift star-forming galaxies comes in several forms, including blueshifts of hundreds of km s$^{-1}$ in interstellar absorption lines relative to galaxy systemic redshifts \\citep{pettini01,shapley03,steidel10}, and the nature of the IGM environments of vigorously star-forming galaxies, in terms of the optical depth and kinematics of the surrounding H I and heavy elements \\citep{adelberger03,adelberger05}. Yet, in spite of evidence for ubiquitous outflows at high redshift, estimates of fundamental physical quantities such as gas column densities and mass outflow rates in the superwinds have remained elusive. While there is much observational evidence for outflows in high-redshift galaxies, both analytic models and hydrodynamical simulations of these systems suggest that, at the same cosmic epochs, they should be rapidly accreting cold gas from the IGM, fueling their active rates of star formation \\citep{birnboimdekel2003,keres2005,keres2009,dekel2009}. Searching for observational signatures of the process of cold gas accretion at high redshift remains an open challenge.  The Ly$\\alpha$ feature is one of the most widely used probes of star formation in both the nearby and very distant universe. While Ly$\\alpha$ photons are initially produced by recombining ionized gas in H II regions, resonant scattering through the interstellar medium (ISM) of galaxies can lead to extreme modulation of the intrinsic emission profile, in both frequency and spatial location. Absorption by dust can completely suppress Ly$\\alpha$ emission, producing a strong absorption profile, even in galaxies with high rates of star formation. It is therefore difficult to determine the intrinsic properties of the gas giving rise to Ly$\\alpha$ emission from the observed profile alone. Only recently has it become possible to make detailed theoretical predictions for the Ly$\\alpha$ profiles emergent from complex systems similar to those observed at high redshift \\citep{ahn02,zheng02,hansenoh06,verhamme06,verhamme08}. These new calculations \\citep[e.g.,][]{verhamme06} employ a Monte Carlo approach to propagate a representative ensemble of Ly$\\alpha$ photons through gas and dust of arbitrary spatial and velocity distribution outputting Ly$\\alpha$ profiles as would be observed. By comparison with Ly$\\alpha$ profiles in actual galaxy spectra, it is possible, in principle, to recover quantities such as the expansion/infall velocity of the outflow/inflow, the column density and velocity dispersion of absorbing gas, and the gas covering fraction. Therefore, modeling of observed Ly$\\alpha$ emission profiles represents an independent method of probing the processes of feedback and accretion at high redshift.  The majority of Ly$\\alpha$ emission profiles at high redshift fall in the category of single-peaked and asymmetric \\citep{shapley03,tapken07}, which is a natural outcome of an expanding medium \\citep{verhamme08}.  However, multiple-peaked Ly$\\alpha$ profiles, seen in a fraction of star-forming UV-selected galaxies at $z\\sim2-3$, offer a particularly detailed perspective on Ly$\\alpha$ escape due to the rich structure of the emergent line.  In principle, the structure of the Ly$\\alpha$ line encodes the velocity field and density distribution of the gas through which it has emerged.  For example, a symmetric double-peaked profile centered on the velocity-field zeropoint is a natural outcome of the radiative transfer of Ly$\\alpha$ photons through a static medium \\citep{osterbrock1962}. Recent advances in modeling the escape of the Ly$\\alpha$ photons from star-forming galaxies at $z\\sim2-3$ have isolated several features of the Ly$\\alpha$ profile that may be expected for specific gas geometries and velocity fields \\citep[e.g.][]{verhamme06,laursen2009a,laursen2009b,barnes2011}. Since Ly$\\alpha$ is the most readily observed high-redshift emission line, there is enormous potential to better understand the structure of the early galaxies by interpreting Ly$\\alpha$ line profiles.  Previous attempts to compare observed Ly$\\alpha$ line morphologies to theoretical predictions have been missing crucial information, weakening the derived constraints on galaxy outflow and inflow properties. In particular, accurate systemic redshift measurements, nebular line widths, and intrinsic ionizing photon fluxes have been absent from most previous comparisons \\citep[but see, e.g.,][]{yang2011,mclinden2011}. These three observables are critical for anchoring the velocity scale of the models, constraining the mass and thermal motions of the gas, and determining the overall normalization for a given model. In this paper, we use new results obtained from H$\\alpha$ and [OIII]$\\lambda$5007 emission lines to provide the critical missing constraints on the observed kinematics of star-forming galaxies at $z\\sim2-3$ with multiple-peaked Ly$\\alpha$ emission.  In addition, we have used our large database of Ly$\\alpha$ emission lines in high-redshift objects to determine the $\\it{frequency}$ of the multiple-peaked systems, providing a global perspective on the potential of using Ly$\\alpha$ morphology to reveal $z\\sim2-3$ gaseous structure.  Our method of target selection from the parent sample of UV-selected galaxies at $z\\sim2-3$ is explained in Section~\\ref{sec:samp}, while the observations and data reduction are described in Section \\ref{sec:obsdata}. The Ly$\\alpha$ velocity profiles are presented in Section \\ref{sec:lyavel} with precise velocity-field zeropoints determined from our H$\\alpha$ (or [OIII]$\\lambda5007$) measurements.  In Section \\ref{sec:physquant} the measured physical quantities for each system are reported.  Section \\ref{sec:model} describes current Ly$\\alpha$ radiative transfer models, in addition to a qualitative comparison between some simple models and our measured Ly$\\alpha$ velocity profiles. Finally, in Section \\ref{sec:conclusions} we summarize our results and discuss how our measurements will be used in future work to accurately model the processes that give rise to the multiple-peaked Ly$\\alpha$ profiles from star-forming galaxies at $z\\sim2-3$.  We assume a flat $\\Lambda$CDM cosmology with $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_{o}=70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusions}  There is strong observational evidence for the ubiquity of large-scale outflows in high-redshift star-forming galaxies. Gauging the overall impact of these large-scale outflows on the evolution of the galaxies sustaining them remains an open challenge. Theoretical considerations suggest that, at the same cosmic epochs, the {\\it inflow} of cold gas is also a very important process in the evolution of star-forming galaxies. Multiple-peaked Ly$\\alpha$ profiles in star-forming UV-selected galaxies at $z\\sim2-3$ offer a unique probe of both outflows and inflows in the early universe. The complex line structure potentially provides constraints on many important gas flow parameters.  Specifically, the locations of peaks with respect to the velocity-field zeropoint, their separations, and flux ratios are all additional pieces of information not seen in single-peaked Ly$\\alpha$ profiles.  Along with observational data such as accurate systemic redshifts, intrinsic nebular line widths, and intrinsic ionizing photon fluxes, one can distinguish among the possible processes underlying these systems.  We have shown that the phenomenon of multiple-peaked profiles appears in a significant fraction ($\\sim$ 30$\\%$) of UV-selected star-forming galaxies at $z\\sim2-3$ with Ly$\\alpha$ emission.  After identifying possible candidates for this study we presented a sample of 18 multiple-peaked objects with observed H$\\alpha$ or [OIII] nebular emission lines, which were used to establish accurate systemic redshift measurements.  The average velocity dispersion, $\\sigma_{v}$, and the mean star-formation rate in our sample are comparable to values measured from the full population of UV-selected star-forming galaxies at $z\\sim2-3$.  At the same time, given our focus on Ly$\\alpha$-emitting galaxies, the average $E(B-V)$ in our sample is bluer than what is typically measured.  Detected interstellar absorption lines for objects in the sample show, on average, a blueshift, which is suggestive of large-scale outflows driven by supernovae or winds of massive stars.  Additionally, the interstellar absorption line widths are measured to be $\\langle \\sigma_{abs} \\rangle\\sim 200$ km s$^{-1}$, which is similar to values found in other studies of star-forming galaxies at $z\\sim2-3$.  A comparison of our interstellar absorption line widths to a sample of star-forming galaxies with single-peaked Ly$\\alpha$ emission show the single-peaked objects to have consistently larger interstellar absorption line widths.  The difference in interstellar absorption line widths may signify that the ranges of gas velocities are different for multiple-peaked compared to single-peaked Ly$\\alpha$ emission objects.  We have qualitatively compared our observed Ly$\\alpha$ profiles with simple models of expanding and infalling gaseous shells and halos, given the attention such models have recently received in the literature \\citep{verhamme06, verhamme08,schaerer2008}. The closest match between these simple models and our data is found for Group I type profiles, which make up 11 out of 18 of the NIRSPEC sample objects, and are characterized by a red-asymmetric double peak straddling the velocity-field zeropoint. Group I type profiles can be produced in an expanding shell model, with a high-column density ($N=2 \\times 10^{20} \\mbox{ cm}^{-2}$), moderate expansion speed ($V_{\\rm exp}=100$~km~s$^{-1}$), and small Doppler parameter, ($b=40$~km~s$^{-1}$). At the same time, the observed low-ionization interstellar absorption line profiles for Group I objects are significantly broader than the features that would arise from an expanding shell with $b=40$~km~s$^{-1}$, and signal a problem with the above interpretation for the Group I Ly$\\alpha$ spectra. Certain aspects of Group II and Group III Ly$\\alpha$ spectra are also reproduced by the simple models considered here, though, in detail, the matches are not perfect. For the remaining handful of objects in our NIRSPEC multiple-peaked sample, the suite of models we've considered cannot be used to explain the observed profiles, both in terms of Ly$\\alpha$ peak locations, relative strengths, and numbers.  It is also of interest to compare Ly$\\alpha$ radiative transfer models with high-redshift data of significantly higher S/N and resolution than the spectra presented here.  \\citet{quider09} present the multiple-peaked Ly$\\alpha$ emission profile of the strongly gravitationally lensed galaxy, ``the Cosmic Horseshoe,\" for which the magnification is estimated to be a factor of $\\sim 24$.  The Cosmic Horseshoe has a very well-defined stellar systemic redshift, from which the Ly$\\alpha$ spectral profile can be calculated in velocity space. While an expanding shell model from \\citet{verhamme06} with $N_{HI}=7 \\times 10^{19} \\mbox{ cm}^{-2}$, $V_{\\rm exp}=300$~km~s$^{-1}$, and $b=40$~km~s$^{-1}$, provides a qualitative match to the two redshifted peaks in the Cosmic Horseshoe Ly$\\alpha$ profile, the small Doppler $b$-parameter in the model is again at odds with the broad low-ionization interstellar absorption profile. Furthermore, in detail, the locations of the Ly$\\alpha$ peaks in the predicted model spectrum do not line up with those of the observations. Both typical and strongly-lensed UV-selected star-forming galaxies with well-determined systemic redshifts and, accordingly, Ly$\\alpha$ profiles calculated robustly in velocity space are now providing important observational constraints on Ly$\\alpha$ radiative transfer models.  These objects highlight the need for considering simultaneously the emergent Ly$\\alpha$ emission profile, and the profile of low-ionization interstellar absorption, which must be matched within a unified framework \\citep{steidel10}.  While we have mainly focused here on the discrepancies between models and data, and have therefore not yet been able to derive the physical parameters of the gaseous flows surrounding UV-selected star-forming galaxies at high redshift, we view this analysis as an important initial step. It will simply not be possible to constrain the gas density and velocity distributions in these circumgalactic flows without the simultaneous establishment of the location of the multiple-peaked Ly$\\alpha$ profile and interstellar absorption features in velocity space. We have presented such measurements here for a sample of objects whose red- and blue-asymmetric multiple-peaked Ly$\\alpha$ profiles at least qualitatively suggest a range of processes including both outflows and infall. Additional critical observables presented here which will be be considered in future modeling are the input velocity distribution of Ly$\\alpha$ photons as traced by the velocity dispersion of either H$\\alpha$ or [OIII] emission, and the degree of dust attenuation as traced by rest-frame UV colors and Ly$\\alpha$/H$\\alpha$ flux ratios.  With a realistic model that successfully matches the combination of intrinsic and emergent Ly$\\alpha$ velocity distributions as well as interstellar absorption profiles, we will obtain the gas density and velocity distribution as a function of radius. These distributions will potentially allow us to constrain the rates at which cold gas mass is flowing in and out of galaxies, which represents one of the most important goals in the study of galaxy formation."
    },
    "1107/1107.0490_arXiv.txt": {
        "abstract": "Scaling relations that link asteroseismic quantities to global stellar properties are important for gaining understanding of the intricate physics that underpins stellar pulsation. The common notion that all stars in an open cluster have essentially the same distance, age, and initial composition, implies that the stellar parameters can be measured to much higher precision than what is usually achievable for single stars.  This makes clusters ideal for exploring the relation between the mode amplitude of solar-like oscillations and the global stellar properties. We have analyzed data obtained with NASA's \\kepler~space telescope to study solar-like oscillations in 100 red giant stars located in either of the three open clusters, NGC~6791, NGC~6819, and NGC~6811. By fitting the measured amplitudes to predictions from simple scaling relations that depend on luminosity, mass, and effective temperature, we find that the data cannot be described by any power of the luminosity-to-mass ratio as previously assumed.  As a result we provide a new improved empirical relation which treats luminosity and mass separately.  This relation turns out to also work remarkably well for main-sequence and subgiant stars. In addition, the measured amplitudes reveal the potential presence of a number of previously unknown unresolved binaries in the red clump in NGC~6791 and NGC~6819, pointing to an interesting new application for asteroseismology as a probe into the formation history of open clusters. ",
        "introduction": "The highly complex processes involved in the excitation and damping of stochastically excited (solar-like) oscillations make estimation of their amplitudes from pulsation modelling particularly challenging \\citep[e.g.~][]{Houdek06,Samadi07}. A scaling relation for the amplitude has therefore been of significant interest since it was first introduced by \\citet{KjeldsenBedding95}.  Their `$L/M$ relation', based on theoretical work by \\citet{DalsgaardFrandsen83} of near main-sequence stellar models, suggested that the amplitude in radial velocity would simply scale as the luminosity-to-mass ratio.  Using observations of stars made in both radial velocity and intensity \\citet{KjeldsenBedding95} also suggested that the amplitude in intensity, $A_\\lambda$, would scale as \\begin{equation} A_\\lambda=\\frac{((L/L_\\odot)/(M/M_\\odot))^s}{\\lambda/500\\mathrm{nm} (T_{\\mathrm{eff}}/5777\\,\\mathrm{K})^r}\\,A_{500\\mathrm{nm},\\odot}, \\label{scaling} \\end{equation} where $s=1$, $\\lambda$ is the central wavelength of the photometric bandpass, and $A_{500\\mathrm{nm},\\odot}$ the observed solar value at 500nm. They found empirically $r=2.0$, which was a slight modification to $r=1.5$ derived if they assumed the stellar oscillations to be purely adiabatic. Subsequent modelling by e.g. \\citet{Houdek99,Samadi07} has lead to variations of the $L/M$ relation where, in essence, different powers of the $L/M$ ratio have been derived ($s=0.7$--$1.3$). Recently, \\citet{Verner11} found $s=0.4$--1.0 depending on \\teff~of a large sample (642) of main-sequence and subgiant stars observed by \\kepler~\\citep{Koch10}.  The existence of solar-like oscillations in red giant stars is now well established observationally, most recently from CoRoT \\citep[e.g.~][]{Ridder09} and \\kepler~\\citep[e.g.~][]{Gilliland10,Bedding10}, as well as theoretically \\citep{Dupret09,Montalban10,Dimauro11}. Despite the significantly different structures of red giants compared to the stars and models on which the $L/M$ relation has been founded, the absence of an alternative has also seen this relation widely used for red giants, including several attempts to determine the best matching exponent, $s$ \\citep[][]{Stello07,Mosser10,Baudin11}. While the majority of results on red giants are on field stars, the recent clear detections in open cluster red giants emerging from \\kepler~\\citep[][Paper I]{Stello10} has opened up the seismic exploration of clusters and the advances that clusters bring to the interpretation of asteroseismic data (\\citealt{Basu11,Hekker10a}; Miglio et al. in prep.; Stello et al. submitted). In particular, stars in an open cluster are thought to share a common distance and initial chemical composition, which allows one to derive the stellar luminosity to much higher precision than for most field stars.  In addition, the common age of red giant stars within each cluster implies that they have practically the same mass, resulting in a relatively low uncertainty on their measured mean mass assuming there is no significant mass loss (Miglio et al. in prep.).  Combined with high quality standard photometry we can therefore obtain more robust predictions of the amplitudes from scaling relations and hence investigate these in ways not possible for the field stars observed by the current space mission CoRoT and \\kepler.  Based on only one month of \\kepler~data of a single cluster, in Paper I we already demonstrated the potential for investigating the $L/M$ relation by taking advantage of the common cluster properties of the stellar sample. We now have \\kepler~time-series photometry that span 10 times longer for stars in three open clusters (NGC~6791, NGC~6819, and NGC~6811), which exhibit distinctly different stellar masses. In this paper we are therefore extending considerably the analysis of the amplitude scaling relation for solar-like oscillations. ",
        "conclusions": "Our analysis of solar-like oscillations in 100 red giant stars in three open clusters revealed that previously adopted scaling relations based on the luminosity-to-mass ratio for predicting amplitudes are not adequate for red giants. We found an empirical scaling relation by fitting the observed amplitudes to a more general form than the previous $L/M$ relation. The result, \\begin{eqnarray} A_\\lambda &\\propto& L^{0.90}/(M^{1.7}T_{\\mathrm{eff}}^{2}) % \\end{eqnarray} \\noindent and \\begin{eqnarray} A_\\mathrm{bol} &\\propto& L^{0.90}/(M^{1.7}T_{\\mathrm{eff}}), \\end{eqnarray} which showed considerable improvement for red giants, turned out to also work remarkably well for main-sequence and subgiant stars.  Interestingly, the lower than expected amplitudes of some red clump stars in NGC~6791 and NGC~6819 revealed that they were likely unresolved binaries, many of which were not known previously. This method for identifying binaries could add interesting new insight to the formation history of these clusters.  In this investigation we ignored any possible effect on amplitude from metallicity differences \\citep{Samadi07} of the three clusters, which have values of [Fe/H]$_{\\mathrm{NGC6791}}\\simeq0.3$ and [Fe/H]$_{\\mathrm{NGC6819}}\\simeq0.1$ \\citep[see~][and reference therein]{Basu11}, while it is unknown for NGC~6811.  To improve on that will require better determination of the cluster metallicities.  In addition, we would need more clusters (with significantly different stellar parameters) to allow the fitting of a rigorous empirical relation including one more free parameter such as metallicity.  With more \\kepler~data in the future, we expect to have cluster stars covering a large range of \\teff, which will include turn-off stars at the end of the main sequence, allowing us to also fit the exponent, $r$, of the \\teff~dependence in the amplitude scaling relation."
    },
    "1107/1107.2389_arXiv.txt": {
        "abstract": "We have discovered a detached pair of white dwarfs (WDs) with a 12.75 min orbital period and a 1,315 \\kms\\ radial velocity amplitude.  We measure the full orbital parameters of the system using its light curve, which shows ellipsoidal variations, Doppler boosting, and primary and secondary eclipses.  The primary is a 0.25 \\msun\\ tidally distorted helium WD, only the second tidally distorted WD known. The unseen secondary is a 0.55 \\msun\\ carbon-oxygen WD.  The two WDs will come into contact in 0.9 Myr due to loss of energy and angular momentum via gravitational wave radiation.  Upon contact the systems may merge yielding a rapidly spinning massive WD, form a stable interacting binary, or possibly explode as an underluminous supernova type Ia.  The system currently has a gravitational wave strain of $10^{-22}$, about 10,000 times larger than the Hulse-Taylor pulsar;  this system would be detected by the proposed LISA gravitational wave mission in the first week of operation.  This system's rapid change in orbital period will provide a fundamental test of general relativity. ",
        "introduction": "In Einstein's general theory of relativity, close pairs of stars produce gravitational waves, ripples in the curvature of space-time \\citep{einstein16,einstein18}.  Observations of the Hulse-Taylor binary pulsar PSR B1913+16 confirm the slow decay of the orbit predicted by Einstein's theory and provide indirect evidence for gravitational waves \\citep{hulse75,weisberg10}. On-going and proposed instruments seek to detect gravitational waves directly \\citep{jafry94,hobbs09,abbott10}.  The strongest known gravitational wave sources are compact binaries containing neutron stars and WDs \\citep[e.g.,][]{nelemans09}. There are presently four known binaries with orbital periods less than 15 minutes -- three AM CVn stars \\citep{haberl95,israel99,warner02} and one low mass X-ray binary \\citep{stella87} -- but all binaries with $<$1 hr orbital periods, with one exception \\citep{kilic11b}, are interacting systems that complicate direct tests of Einstein's theory.  \\citet{marsh05} show that the expected period change due to gravitational wave radiation can be modified by both induction-driven angular momentum interchange in magnetic WD systems and by mass-transfer in accreting systems.  Moreover, there are different explanations for both the nature of the shortest-period AM CVn systems and their observed period changes \\citep{strohmayer04,dantona06,deloye06,roelofs10}.  Here, we report the discovery of a detached, 765 sec orbital period binary, SDSS J065133.33+284423.3 (hereafter J0651).  The discovery comes from our on-going ELM Survey, a targeted spectroscopic survey for extremely low mass $<$0.25\\msun\\ WDs \\citep{brown10c,kilic10}.  The ELM Survey demonstrates that extremely low mass WDs are found in compact binaries, and that the majority of these binaries have $<$10 Gyr merger times due to gravitational wave radiation \\citep{kilic11a}. Depending on the poorly constrained physics of the merger process, these merging WD systems may produce single helium-rich sdO stars, stable mass-transfer AM CVn binaries, or possibly underluminous supernova \\citep[as discussed in][]{kilic10}. The observed merger rate of the low mass WD binaries is about 1\\% of the Type Ia supernovae rate -- comparable to the rate of underluminous supernovae \\citep{brown11a}.  The ELM Survey is also responsible for finding the first tidally distorted WD \\citep{kilic11b}.  J0651 is the most extreme system discovered to date and shows that the ELM Survey may open a new window on gravitational wave sources.  In J0651, the orientation of the binary allows us to observe eclipses of each star by the other, leading to accurate measurement of the orbital parameters, masses, and WD radii.  There is no evidence for mass transfer.  These results suggest that J0651 is the cleanest known strong gravitational wave source, with an orbital decay time of less than 1 Myr.  Future observations will allow us to measure the change in the orbital period.  We hope to one day compare this change with direct measurements of gravitational waves and provide an unprecedented test of general relativity.  ~  \\begin{figure}\t\t% \\plotone{f1.eps} \\caption{ \\label{fig:spec} The composite observed spectrum of J0651, constructed by shifting the individual spectra to rest-frame and then summing them together to create an effective 60.5 min exposure, compared to the best-fit stellar atmosphere (solid red line)  for a 0.25 \\msun\\ WD.  The broad hydrogen Balmer absorption lines visually indicate J0651 is an unusually low mass WD.} \\end{figure} ",
        "conclusions": "The eclipsing detached WD binary system J0651 presents us with a remarkable laboratory.  We can use its eclipses to measure WD masses and radii for detailed tests of WD models.  We can use its changing binary orbital period and gravitational wave strain to test for the existence of gravitational waves predicted by general relativity.  We can use its very existence to constrain the space density and merger rate of low mass WD binaries and links to underluminous supernovae.  In the future, we plan to use multi-passband photometry to directly measure the nature of the secondary WD and to detect the change in orbital period predicted by general relativity."
    },
    "1107/1107.1033_arXiv.txt": {
        "abstract": "We examine the dark energy and matter densities allowed by precision measurements of distances out to various redshifts, in the presence of spatial curvature and (near) arbitrary behavior of the dark energy equation of state. Degeneracies among the parameters permit a remarkably large variation in their values when using only distance measurements of the late time universe and making no assumptions about the dark energy or curvature.  Going beyond distance measurements to a lower limit on the growth of structure bounds the allowed region significantly but still leaves considerable freedom to match a flat $\\Lambda$CDM model with dark energy very different from a cosmological constant.   The combination of distances with Hubble parameter, gravitational lensing or other large scale structure data is essential to determining robustly the cosmological model. ",
        "introduction": "A central goal of modern cosmology is to reveal in detail the energy budget of the universe.  In addition to matter (baryonic and dark matter) and small contributions by radiation and neutrinos, there is an (effective) dark energy associated with the accelerated expansion and possibly an (effective) curvature energy associated with deviation from spatial flatness.  Cosmological observations, especially over the past decade, have made great strides in constraining the energy density fractions in each of these, but generally assuming specific behaviors for the dark energy.  Since dark energy is an almost total mystery, it behooves us to reexamine the issue and ask to what extent the $\\Lambda$CDM concordance model of $\\sim28\\%$ matter density and $\\sim72\\%$ cosmological constant energy density is necessarily close to the true energy budget.  Because all the energy densities enter into the Hubble expansion rate, which then determines the distance-redshift relation, degeneracies exist between the components such that more of one can compensate for less of another.  Since they evolve differently with redshift, however, each characterized by their own equation of state parameter (0 for matter, $-1/3$ for curvature, $w(z)$ for dark energy), one expects that observations over a sufficiently wide redshift range give leverage to break the degeneracies.  This has been explored for restricted scenarios of matter and dark energy densities (e.g.\\ \\cite{steinhardt,Shafieloo05,Shafieloo07,om}) and curvature and dark energy densities (e.g.\\ \\cite{lincurv,bassett}), and nonparametrically from the observations through redshifts bins of dark energy (e.g.\\ \\cite{union21}).  Perhaps closest in philosophy to our approach are the works of \\cite{mhh1,mhh2,mhh3}, which look at how the diversity of models translate to dispersion in observables, while we explore the converse of how a tight observable relation can arise from a wide range of models.  In this paper we investigate the freedom around the concordance model caused by degeneracies when we allow for matter, curvature, and dark energy with no a priori restriction on its equation of state.  The philosophy is to use as direct measurements as possible without making assumptions about the dark energy.  Therefore we restrict ourselves to late time observations since we have no knowledge of dark energy behavior at early times, e.g.\\ is there early dark energy affecting the cosmic microwave background (CMB).  We use purely geometric distance measurements, examining how the degeneracies are broken as the data quality and redshift range improves.  That is, suppose we have even exact agreement with a particular flat $\\Lambda$CDM model in the distance-redshift relation out to some redshift $z$; how close to the concordance model in the matter density--dark energy density plane does this restrict the cosmological model when allowing for curvature and arbitrary $w(z)$ behavior? We should note that in this paper we are considering the degeneracy between the fundamental cosmological quantities namely matter density, curvature and the effective equation of state of dark energy. There are also another sort of degeneracies between different cosmological models which are very different by nature but they result to similar effective equation of state of dark energy~\\cite{SahniShtanov1,SahniShtanov2}.  Section \\ref{sec:distance} lays out the methodology and explains the constraints imposed by various levels of conditions imposed by consistency and by the observations.  The influence of measurements apart from distance, such as age of the universe and the linear growth factor of cosmic structure, are also addressed.  In Section~\\ref{sec:degen} we quantify the constraints in the matter--dark energy density ($\\om$--$\\ode$) plane and Section~\\ref{sec:concl} summarizes the results about how well we actually know our cosmological model. ",
        "conclusions": "}  We have quantified the cosmographic degeneracy present when matter, spatial curvature, and unrestricted dark energy models can contribute to the distance-redshift relation.  Even when perfectly matching the distances out to $\\zmax=1.5$ for a flat $\\Lambda$CDM model with a given matter density, a substantial region of the density parameter space remains degenerate with the true model.  This implies that we {\\it cannot\\/} assume that the cosmological constant describes the dark energy through such distance measurements alone.  Imposing other low redshift constraints, such as basic consistency conditions on the radius of curvature of closed universes and positivity of the dark energy density, and observational criteria such as a minimum age of the universe and a simple lower bound on the total growth factor for large scale structure, still leaves considerable freedom for the curvature and dark energy contributions.  One would zero in on a $\\Lambda$CDM model only when restricting the form of the dark energy evolution or bringing in early universe constraints (e.g.\\ CMB and more detailed large scale structure characteristics) -- assuming dark energy has negligible contribution at high redshift.  Distances involving curvature differently, e.g.\\ parallax \\cite{weinberg} or distance intervals, including direct measurement of the Hubble parameter or gravitational lensing, of galaxies or the CMB, may offer another path (see e.g.\\ \\cite{bernstein,knox,kazin,suyu}).  We emphasize that we have made essentially no assumptions about the dark energy behavior, allowing its equation of state parameter $w$ to range from $-\\infty$ to $+\\infty$ (with the exception of restricting $w\\le1$ above $\\zmax$ when using the minimum growth condition).  A constraint on $w$ to lie within $[-1,+1]$, say, the values for a canonical scalar field, would limit the allowed region to a much smaller area around the true values in the density space. When the assumed $\\Omega_{m}$ is larger than the actual matter density, at some redshifts we need a very large negative equation of state of dark energy to suppress the contribution of dark energy in the total energy density of the universe.  Conversely, when the assumed $\\Omega_m$ is smaller than the actual matter density we need a dark energy with equation of state of greater than $-1$ and in some cases even greater than $1$ to compensate the lack of contribution from the matter part.  Thus, imposing the limit $w < +1$, say, rules out an area with low matter and dark energy densities. A constraint on spatial curvature (while still allowing for arbitrary dark energy behavior) would cut a diagonal swath in the density plane parallel to the flatness line, considerably restricting the allowed region.  It is interesting to see how our ignorance of the nature of dark energy and the geometric curvature of space diffuse the strength of evidence for the cosmological constant model from distance measurements.  The true universe may be much more complicated, and yet perfectly consistent with cosmography, than this highly restricted model.  Combining distance measurements with gravitational lensing data and mapping of large scale structure will greatly reduce the degeneracies exhibited and such a suite of future observations offers true hope to understand our universe in a less model dependent manner."
    },
    "1107/1107.2051.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {Ceres is the most massive body of the asteroid belt and contains about 25 wt.\\% (weight percent) of water. Understanding its thermal evolution and assessing its current state are major goals of the \\textit{Dawn} Mission. Constraints on internal structure can be inferred from various observations. Especially, detailed knowledge of the rotational motion can help constrain the mass distribution inside the body, which in turn can lead to information on its geophysical history.} % aims heading (mandatory) {We investigate the signature of the interior on the rotational motion of Ceres and discuss possible future measurements performed by the spacecraft \\textit{Dawn} that will help to constrain Ceres' internal structure. } % methods heading (mandatory) {We compute the polar motion, precession-nutation, and length-of-day variations. We estimate the amplitudes of the rigid and non-rigid response for these various motions for models of Ceres interior constrained by recent shape data and surface properties. } % results heading (mandatory) {As a general result, the amplitudes of oscillations in the rotation appear to be small, and their determination from spaceborne techniques will be challenging.  For example, the amplitudes of the semi-annual and annual nutations are around $\\sim 364$ and $\\sim 140$ milli-arcseconds, and they show little variation within the parametric space of interior models envisioned for Ceres. This, combined with the very long-period of the precession motion, requires very precise measurements. We also estimate the timescale for Ceres' orientation to relax to a generalized Cassini State, and we find that the tidal dissipation within that object was \\tnr{probably} too small to drive any significant damping of its obliquity since formation. However, combining the shape and gravity observations by \\textit{Dawn} offers the prospect to identify departures of non-hydrostaticity at the global and regional scale, which will be instrumental in constraining Ceres' past and current thermal state. We also discuss the existence of a possible Chandler mode in the rotational motion of Ceres, whose potential excitation by endogenic and/or exogenic processes may help detect the presence of liquid reservoirs within the asteroid.  } % conclusions heading (optional), leave it empty if necessary ",
        "introduction": "The largest asteroid, Ceres, will be rendez-vous'd by the {\\it{Dawn}} spacecraft in 2015. During the nominal mission, the asteroid will be mapped with high-resolution imaging in the visual and infra-red wavelengths (VIR instrument), and its gravity field will be determined to the tenth degree and order (McCord et al. 2011). A major goal of the {\\it{Dawn}} Mission is to constrain the internal structure of Ceres, and especially quantify the extent of differentiation, and thus of internal evolution, achieved by this dwarf planet (Russell \\etal 2007; McCord \\etal 2011).  While the primary objective of high-resolution camera VIR is to map the surface composition and surface feature, an interesting contribution of these observations will be to characterize Ceres' rotational motion. Rotation properties have proved to bring important constraints on the interiors of planetary bodies, (e.g., Mathews \\etal 2002; Koot \\etal 2010; Williams \\etal 2001; Margot \\etal 2007; Dehant \\etal 2009) because they depend on mass distribution and viscoelastic properties. Indeed, the departures from a uniform rotation and changes in the orientation of a body are responses to an external forcing such as the gravitational force of another celestial body of the solar system or of the Sun. These responses depend on the structure and composition of the interior. In particular the possible presence of liquid layers inside a body and the elastic or inelastic properties of the solid parts drive the object's response to the external forcing. In particular when there are liquid layers within a body, such as Mercury or the icy satellites, the librational response of the solid part is amplified. A similar effect is expected for the nutation when considering rapid orientation changes such as for the Earth and Mars (resonance to the Free Core Nutation). The material composing Ceres and its physical properties thus determine the rotational behavior of Ceres. The observation of the response of Ceres to any forcing may thus provide information on its interior.  The purpose of this paper is to investigate the temporal variations of Ceres' rotation in order to identify potential observations to be performed by the \\textit{Dawn} Mission. So far, observations of Ceres' surface have been obtained with the \\textit{Hubble Space Telescope} ({\\it{HST}}), adaptive optics (Keck Telescope), and occultations. However, it has proved difficult to track surface landmarks with enough accuracy to constrain the polar orientation and rotation of the asteroid (Thomas \\etal 2005; Carry \\etal 2008; Drummond and Christou 2008). Due to the non-spherical shape of Ceres related to the centrifugal potential, the gravitational potential of the Sun exerts a non-zero torque on Ceres dynamical figure. Consequently, Ceres' axis of figure is expected to exhibit a precessional motion around its normal to the orbital plane and the periodic part of the torque generates a nutational motion along the precessional cone. In addition, in the reference frame of Ceres, the spin axis describes a polar motion or wobble around the figure axis, and, finally, tidal deformations arising from the Sun induce perturbations in its rotational velocity, leading to length-of-day variations (l.o.d.).  The first part of this article describes the shape and interior structure parameters used for modeling Ceres. Then, we review the observed pole position and discuss briefly its secular evolution. In Section 4, we introduce the main equations used to determine the rotational parameters of the body. The geophysical and rotational models are then combined to compute the rotational motion of Ceres as a function of polar motion and precession-nutation (Section~\\ref{sec:rigide}). Based on these results, we discuss the prospects for characterizing Ceres' rotation with the {\\it{Dawn}} Mission, and the constraints these observations will provide on Ceres' interior (Section~\\ref{sec:geocons}). ",
        "conclusions": "We have characterized the main components of Ceres' rotation and quantified them assuming Ceres is differentiated into a rocky core and icy shell. First, our modeling predicts that Ceres' obliquity is not constrained by the dissipative history of the asteroid. However, multiple determinations of Ceres' pole agree that its obliquity lies between 0 and 4 deg. The lower bound suggests that Ceres could be relaxed to a generalized Cassini state. However, considering the very long damping timescale, such a situation is unlikely. \\tjc{This uncertainty will be resolved by the {\\it{Dawn}} }. In any case, important constraints can also be inferred from combining shape and gravity data. These will yield independent determinations of the secular Love number that will be used to constrain Ceres' hydrostatic state, from which the mean moment of inertia of the asteroid can be inferred.  %\\tnr{The determination of the obliquity of Ceres based on observational data disagree with the assumption of Bills and Nimmo (2010) of 12 deg.}  \\tjc{Faut-il repeter cela? Then, for the stratified, solid model considered in this study, we established upper bounds on the rigid and non-rigid components of the nutations, polar motion and l.o.d. These appear too small to be inferred from space measurement techniques. Then we identified that a detectable perturbation of Ceres' spin state (wobble) may be the signature of a Chandler mode. This mode would have to be excited by recent large impacts or currents in local liquid reservoirs at depth in order to yield a sizeable signature. This aspect needs to be quantified in detail as it offers the prospect to constrain Ceres' thermal state and geophysical evolution.   %KEEP OR REMOVE? The {\\it{Dawn}} Mission will rendez-vous Ceres in 2015 and will nominally spend at least six months in orbit around the asteroid. The {\\it{Dawn}} spacecraft carries a payload composed of several instruments that will help to constrain the physical properties and chemistry of Ceres (Russell \\etal 2006). Of specific relevance to the present study, the camera, visible Infra-Red camera and the radio science subsystems will yield refined estimates of the global shape and the first gravity data for that object. Topography and gravity mapping can be combined to constrain the presence of non-hydrostatic anomalies. Surface mapping will also be used to refine the pole orientation."
    },
    "1107/1107.3429.txt": {
        "abstract": "\\noindent Effects of stellar rotation on adiabatic oscillation frequencies of $\\beta$ Cephei star are discussed. Methods to evaluate them are briefly described and some of the main results  for four specific stars are presented. ",
        "introduction": "Main sequence (MS) massive stars are usually fast rotators and their fast rotation affects their internal structure  as well as their evolution. The issue which is adressed here is what information can we obtain- about rotation - from the oscillations of these  massive, main sequence stars ?   The following seismic diagnostics for rotation using non axisymmetric modes will be discussed: {\\it 1) rotational splittings}  as  direct   probes of     the rotation profile. More precisely, we study the  effects of cubic order in the rotation rate compared to effects of a  latitudinal dependence of the rotation on the splittings; {\\it 2) splitting asymmetries}  as a probe for centrifugal distorsion. The case of  {\\it 3) axisymmetric modes}  as   indirect probes   of rotation throughout  effect of rotationally induced mixing on the structure will also be considered.    Results discussed here are obtained with perturbation methods. For nonperturbative methods and results, we refer the reader to Ligni\\`eres {\\sl et al.} (2006),  Reese {\\sl et al.} (2009) and references  therein.   The paper is organized as follows: in Sect.2, properties of pulsating B  stars are recalled with emphasis on the uncertainties of their physical description that can be addressed by seismic analyses. Sect. 3. recalls the  theoretical framework for seismic analyses of relevance here. In Sect.4, seismic analyses of four  $\\beta$ Cep  are presented. In Sect.5 a theoretical  study compares the modifications of the rotation splittings due either to latitudinal dependence of the rotation rate, $\\Omega$, or to cubic order ($O(\\Omega^3$)) frequency corrections. Some conclusions are given in Sect.6. ",
        "conclusions": "We have seen along this review that several efficient seismic tools  can be designed   to obtain valuable information on the internal structure and dynamics of  main sequence massive  stars which  oscillate with a few identified modes. Identification of the detected  modes requires a high signal to noise which is made available due to the large amplitudes of these opacity-driven modes. On the other hand, these stars oscillate with low  frequencies lying near/in the dense part of the spectrum where  p modes, mixed modes and g modes can be encountered. While this is a great advantage in order  to probe the inner layers of the star, resolution and precise measurement of quite close  frequencies in a Fourier spectrum requires  very  long time series. This explains the yet still small number of $\\beta$ Cephei stars for which a successful seismic analysis has been obtained, despite the appealing prospects that a better knowledge of their structure bring up valuable constrains on  their still poorly understood life end as supernovae. It is expected that the space  experiments CoRoT (Michel {\\sl et al.} 2008) and Kepler (Christensen-Dalsgaard \\etal, 2008) will  increase the number of O-B  stars for which fruitful seismic analyses can be carried out as well as possibly enlarge the sample to fast rotators.  Mode identification can be at first difficult to perform for fast rotators but some of these fast rotating stars might also  show  oscillations  of solar-like type which characteristics could help the mode identification. This interesting perspective has recently emerged  with the discovery of the first chimera star    with the CoRoT mission (Belkacem {\\sl et al.}, 2009)."
    },
    "1107/1107.1049_arXiv.txt": {
        "abstract": "{Dust attenuation curves in external galaxies are  useful to study their dust  properties as well as to interpret their intrinsic spectral energy distributions. These functions are not very well known in the UV range either at low or high redshift. In particular the presence or absence of a UV bump at 2175 $\\AA$ remains an open issue which has consequences on the interpretation of broad band colours of galaxies involving the UV range.} {We want to study  the dust attenuation curve in the UV range at $z >1$  where the UV is redshifted into the visible, and with $Herschel$ data to constrain dust emission and a global dust attenuation. In particular we will search for the presence of a UV bump  and the related implications on dust attenuation determinations.} {We use deep photometric data of the Chandra Deep Field South obtained  with  intermediate and broad band  filters  by the MUSYC project  to sample the UV rest-frame of galaxies with  $1<z <2$. {\\it Herschel}/PACS  and  {\\it Spitzer}/MIPS data are used to measure the dust emission. 30 galaxies were  selected  with  high S/N in all  bands. Their SEDs from the UV to the far-IR are fitted using the  CIGALE code  and the characteristics of the  dust attenuation curve are obtained as Bayesian outputs of the SED fitting process. } { The  mean dust attenuation curve we derive exhibits  a significant UV bump  at 2175 $\\AA$ whose amplitude  corresponds  to 35$\\%$ (76 $\\%$)  that of the Milky Way (Large Magellanic Cloud: LMC2 supershell) extinction curve.  An analytical expression of the average attenuation curve ($A(\\lambda)/A_{\\rm V}$)  is given, it is found slightly steeper than  the  Calzetti et al.  one, although at a 1$\\sigma$ level.  Our  galaxy sample is used to study  the derivation of the  slopes of the UV continuum  from broad band colours, including the rest-frame GALEX $FUV-NUV$ colour. Systematic errors induced by the presence of the bump are quantified. We compare dust attenuation factors measured with  CIGALE   to the slope of the UV continuum and find that there is a large scatter   around the relation  valid for local starbursts ($\\sim 0.7$ mag). The uncertainties on the determination of the UV slope   lead to an extra systematic  error of the order of 0.3 to 0.7 mag on dust attenuation when a filter overlaps the UV bump.} ",
        "introduction": "Stellar light in galaxies is absorbed and scattered by the interstellar medium, because of the presence of dust grains. On a galaxy scale, the process is usually quantified by the introduction of an average attenuation function which reflects the combined effects of  absorption and scattering in a complex geometry. \\\\ The most important feature in the extinction curves of the Milky Way (MW) and Large Magellanic Cloud (LMC) is the so called UV bump, a strong absorption feature at 2175 $\\AA$ which is not observed  in the Small Magellanic Cloud (SMC) bar, although a line of sight through the SMC wing exhibits an extinction curve with a strong UV bump \\citep{li02,cartledge05}. The exact origin of this feature is still under discussion \\citep[e.g.][and references therein]{draine89,xiang11}. When dealing with average attenuation curves, the search for a bump is  complicated by the effects of scattering or  geometry \\citep{inoue06,noll07,panuzzo07} which may affect the amplitude of the resulting bump in an average attenuation curve even when assuming  a local Milky Way like extinction curve. \\\\ The best way to study UV bumps  in external galaxies is to perform  UV rest-frame spectroscopy: \\citet{calzetti94} deduced the net attenuation curves of local starburst galaxies  by comparing the UV spectra of  dusty and dust-free spectra of  local starbursts: they did not find any  bump.  Conversely,  \\citet{noll05,noll09a} observed a significant 2175 $\\AA$ absorption feature in  the spectra of  star forming galaxies at $z\\sim 2$ consistent with an extinction curve intermediate between the SMC and LMC ones.  \\\\ The 2175 $\\AA$ feature has also been detected in the spectral energy distribution of some distant galaxies  hosting    gamma-ray bursts \\citep{liang09,liang10}, the extinction curves are found to display a large diversity of shapes  \\citep[][and references therein]{liang09,liang10,zafar11}. We  also refer to \\citet{xiang11} for an extensive review on the detection of the UV bump in Galactic and extragalactic interstellar medium.\\\\ Broad band colours are also used to search for evidence of a UV bump. From an analysis of the observed B-R colours of galaxies at $z\\sim 1$  \\citet{conroy10a}  disfavours  the presence of UV bumps as strong as that in the Milky Way whereas \\citet{burgarella05} analysed the broad bands SEDs of nearby galaxies observed by GALEX and IRAS and conclude to the presence of a bump in the attenuation curve of these galaxies. In all cases, the use of broad band colours which overlap the bump makes difficult to disentangle the effects of a dust extinction curve with or without a bump, of  dust/stars geometry and of various star formation histories \\citep{inoue06,panuzzo07,johnson07a,johnson07b,buat11} leading to contradicting conclusions on the presence of a UV bump in  nearby galaxies observed by GALEX and SDSS \\citep{conroy10b,johnson07a,johnson07b}. \\\\ The presence and characteristics of the UV bump, and far UV rise in the extinction and attenuation curves can provide information on the interstellar medium properties \\citep{noll09a}. Beyond the studies of the interstellar medium, an  reliable attenuation curve  in the UV range is very useful to derive star formation rates in galaxies since the UV emission is commonly used as a star formation tracer.  For example, \\citet{wije10} compared star formation rates deduced from luminosities measured in  the GALEX bands and the H$\\alpha$ and [OII] emission lines and concluded that they must remove the 2175 $\\AA$ feature from the attenuation curve to obtain the best agreement between the different SFR estimators. The   attenuation curve of \\citet{calzetti00} is commonly used, especially in high-z studies and does not include a bump: if  such a bump is present, omitting it leads to incorrect dust attenuation corrections \\citep{ilbert09}. \\\\ The global amount of dust attenuation is robustly estimated by comparing dust and stellar emission, through the $L_{\\rm IR}/L_{\\rm UV}$ ratio. When IR data are not available the shape of the UV continuum ($< 3000$  $\\AA$) has been proposed  as a proxy to measure the amount of attenuation. \\citet{calzetti94} showed that  the UV continuum of local starburst galaxies can be  fit by  a power law  ($ f_\\lambda ({\\rm erg ~cm^{-2} s^{-1}  \\AA^{-1}}) \\propto \\lambda^\\beta$)  for $\\lambda > 1200 $  $\\AA$. They calculated $\\beta$ by defining 10 windows in the IUE spectra of their galaxies  avoiding the 2175 $\\AA$ feature.  \\citet{Meurer99} found a relation between  $\\beta$ and  the amount of dust attenuation measured by $L_{\\rm IR}/L_{\\rm UV}$.  This  local starburst relation is widely used to estimate dust attenuation in high redshift galaxies using a broad band  colour  in the UV rest-frame  of the galaxies as a proxy of  $\\beta$ \\citep[e.g.][]{burgarella07,elbaz07,daddi07,reddy08}. If a bump is present in the dust attenuation curve, one must be cautious by choosing broad band filters which do not overlap the bump in order to estimate reliable slopes outside the wavelength range affected by the bump, as was done to derive the \\citet{Meurer99} relation. This has been  illustrated by studies of the local universe based on GALEX FUV and NUV data: the NUV band largely overlaps the UV bump and the FUV-NUV colour is found to be sensitive to the presence or absence of this feature as well as to the recent star formation history \\citep{panuzzo07,inoue06}.\\\\ In this paper we take advantage of the availability of high quality deep optical data obtained through intermediate-band filters in the Chandra Deep Field South (CDFS) to   sample the rest-frame UV of galaxies at high redshift. These data are combined with the  $Herschel$ observations of the CDFS  field as part of the open time key program GOODS-$Herschel$  project (P.I. D. Elbaz).  The galaxies detected  in the  far-IR   are expected to experience a substantial extinction, which    helps to detect the imprint of the bump in their SED. The combination of UV and IR data will allow us to compare the properties of the UV continuum to the dust attenuation at work in these objects. In a similar vein, \\citet{ilbert09} used broad and intermediate band filters to derive photometric redshifts in the COSMOS field and found that a UV bump is necessary to explain the UV rest-frame fluxes of some starburst galaxies but they did not discuss the amplitude of the bump and  dust attenuation in  these galaxies.\\\\ The construction of the sample is explained in section 2. We want to use only  sources detected with high signal to noise ratio (SNR) in a large number of filters which sample well the rest-frame UV wavelength region affected by the bump.  This criterium leads to a redshift selection between 0.95 and 2.2. The dust attenuation curve as well as the amount of dust obscuration is  constrained by fitting the whole SEDs of the galaxies with the CIGALE code, which is described in section 3. Section 4 is devoted to the results about the presence of the UV bump and section 5 to the consequences on the determination of the slope of the UV continuum and dust attenuation in galaxies. The conclusions are presented in section 6. ",
        "conclusions": "We have identified the presence of a UV bump in  a sample of  galaxies of the CDFS at a redshifts between $0.9 < z < 2.2$, selected to be observed in intermediate and broad band optical filters as well as at mid-and far-IR wavelengths  by {\\it Spitzer} and $Herschel$. This study demonstrates the capabilities of intermediate-band photometry to give details on SED of galaxies, and to derive physical properties. \\\\ The average dust attenuation curve we deduce is well described by a modified Calzetti et al. law slightly steeper than the original one and with a UV bump at 2175 $\\AA$ whose amplitude is $\\sim$ 35$\\%$ of the MW one. We propose an  analytical expression of the average attenuation curve which can  be used to correct the SEDs of  galaxies for  dust attenuation.The moderate amplitude of the bump together with the substantial slope of our average attenuation curve in the UV argue for a deficit of UV bump carriers in our sample galaxies as compared to   the MW.\\\\ Our sample contains seven X-ray galaxies. Five of them are reasonably fitted: a steep attenuation curve is favoured excluding a gray extinction. The sample is too small to draw any  firm conclusion about the presence of a UV bump. In any case the amplitude of the bump, if any, is smaller than that found for non X-ray galaxies.\\\\ The presence of a bump in the mean  dust attenuation curve has implications for the derivation of  the slope of the UV (rest-frame) continuum of galaxies from broad band data alone.  The power law modeling of the UV spectral distribution of our galaxies can be extended up to 3000 $\\AA$ but the bump area must be avoided. When a broad band filter overlaps the bump, despite its moderate amplitude, the error on the determination of the slope may reach $\\sim$0.5.  The use of GALEX  filters leads to an underestimation  of the slope  of only  0.2, even if the NUV filter overlaps the bump. This is likely to be due to the very large bandpass of this filter. However, the proximity of the two central wavelengths of the GALEX filters implies a large uncertainty in the determination of the slope.  It is stressed that a larger wavelength baseline would give a  better measure. \\\\ Dust attenuation estimated with  UV slopes, even when the latter are reliable,  remains uncertain, with an RMS error of 0.7 mag. \\\\ \\onlfig{1}{ \\begin{figure} \\centering \\includegraphics[width=5.8cm,angle=-90]{N1-2sed.pdf} \\includegraphics[width=8cm]{N1-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N2-2sed.pdf} \\includegraphics[width=8cm]{N2-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N3-2sed.pdf} \\includegraphics[width=8cm]{N3-im.pdf}  \\caption{Spectral energy distributions for the full sample and HST images obtained from the GOODS cutout service V0.2. The best fit is represented with a solid line. The rest-frame wavelength is plotted on the x-axis and the fluxes in Jy on the y-axis.  } \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N4-2sed.pdf} \\includegraphics[width=8cm]{N4-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N5-2sed.pdf} \\includegraphics[width=8cm]{N5-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N6-2sed.pdf} \\includegraphics[width=8cm]{N6-im.pdf} \\caption{ cont'd} \\label{fig:continued:second} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N7-2sed.pdf} \\includegraphics[width=8cm]{N7-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N8-2sed.pdf} \\includegraphics[width=8cm]{N8-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N9-2sed.pdf} \\includegraphics[width=8cm]{N9-im.pdf} \\caption{ cont'd} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N10-2sed.pdf} \\includegraphics[width=8cm]{N10-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N11-2sed.pdf} \\includegraphics[width=8cm]{N11-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N12-2sed.pdf} \\includegraphics[width=8cm]{N12-im.pdf} \\caption{ cont'd} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N13-2sed.pdf} \\includegraphics[width=8cm]{N13-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N14-2sed.pdf} \\includegraphics[width=8cm]{N14-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N15-2sed.pdf} \\includegraphics[width=8cm]{N15-im.pdf} \\caption{ cont'd} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N16-2sed.pdf} \\includegraphics[width=8cm]{N16-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N17-2sed.pdf} \\includegraphics[width=8cm]{N17-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N18-2sed.pdf} \\includegraphics[width=8cm]{N18-im.pdf} \\caption{ cont'd} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N19-2sed.pdf} \\includegraphics[width=8cm]{N19-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N20-2sed.pdf} \\includegraphics[width=8cm]{N20-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N21-2sed.pdf} \\includegraphics[width=8cm]{N21-im.pdf} \\caption{ cont'd} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N22-2sed.pdf} \\includegraphics[width=8cm]{N22-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N23-2sed.pdf} \\includegraphics[width=8cm]{N23-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N24-2sed.pdf} \\includegraphics[width=8cm]{N24-im.pdf} \\caption{ cont'd} \\end{figure} \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N25-2sed.pdf} \\includegraphics[width=8cm]{N25-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N26-2sed.pdf} \\includegraphics[width=8cm]{N26-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N27-2sed.pdf} \\includegraphics[width=8cm]{N27-im.pdf} \\caption{ cont'd} \\end{figure}  \\begin{figure} \\ContinuedFloat \\centering \\includegraphics[width=5.8cm,angle=-90]{N28-2sed.pdf} \\includegraphics[width=8cm]{N28-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N29-2sed.pdf} \\includegraphics[width=8cm]{N29-im.pdf} \\includegraphics[width=5.8cm,angle=-90]{N30-2sed.pdf} \\includegraphics[width=8cm]{N30-im.pdf} \\caption{ cont'd} \\end{figure}  }"
    },
    "1107/1107.4988_arXiv.txt": {
        "abstract": "In this paper we discuss the recently discovered radio transient in the nuclear region of M\\,82. It has been suggested that this source is an X-ray binary, which, given the radio flux density, would require an X-ray luminosity, $L_{\\rm X}\\,\\sim6\\times10^{42}$\\,erg\\,s$^{-1}$ if it were a stellar mass black hole that followed established empirical relations for X-ray binaries. The source is not detected in the analysis of the X-ray archival data. Using a 99\\,\\% confidence level upper limit we find that $L_{\\rm X} \\leq 1.8 \\times 10^{37}$\\,erg\\,s$^{-1}$ and $1.5 \\times 10^{37}$\\,erg\\,s$^{-1}$, using powerlaw and disk blackbody models respectively. The source is thus unlikely to be a traditional microquasar, but could be a system similar to SS\\,433, a Galactic microquasar with a high ratio of radio to X-ray luminosity. ",
        "introduction": "\\par M\\,82 is prolific star-forming galaxy at a distance of 3.6\\,Mpc \\citep{1994ApJ...427..628F}. Its proximity therefore makes it an ideal place to study star formation. To this end, M\\,82 has been well monitored by various observatories at several wavelengths for decades \\citep[e.g.][]{1980ApJ...235..392T, 1981ApJ...246..751K, 1988Natur.334...43B,1994MNRAS.266..455M, 1997A&A...320..378S, 2001A&A...365..571W, 2008MNRAS.391.1384F}. In particular, radio observations have revealed approximately 60 compact radio sources in the central region of M\\,82 \\citep{2002MNRAS.334..912M}. A quarter of these sources are of unknown origin. In addition to compact sources, frequent radio observations have also found transient sources of an undetermined nature in M\\,82 \\citep{1985Sci...227...28K, 1994MNRAS.266..455M} and monitored the evolution of radio supernovae \\citep[e.g.][]{1994MNRAS.266..455M, 2006MNRAS.369.1221B, 2008MNRAS.391.1384F}.  \\par The next generation of radio telescopes should reveal many more such transient sources. These observatories will have the ability to detect a 0.1\\,mJy source within minutes. We will then more easily be able to observe transient populations not only in star-forming galaxies like M\\,82 where these systems are abundant, but also in galaxies with low star-formation rates like ellipticals. By attempting to determine the nature of the unknown transients such as those in M\\,82, we will make the task of studying the large number of transients expected to be observed in the future much easier.  \\par In this paper we investigate further the most recently discovered radio transient of unknown origin in M\\,82 and show that this source is unlikely to be a normal microquasar. We discuss the possibility that the M\\,82 source could be an extragalactic analogue of SS\\,433, a very unusual Galactic microquasar \\citep[][and references therein]{1984ARA&A..22..507M}. Both sources have low X-ray and high radio luminosities. In addition, SS\\,433 possesses precessing, helical jets. SS\\,433 also has highly blue- and redshifted optical emission lines associated with it \\citep{1979ApJ...230L..41M}. If the M\\,82 transient is an extragalactic SS\\,433 like microquasar, it will be only the second such source, after the microquasar S\\,26 in NGC\\,7793, to be discovered \\citep{2010Natur.466..209P}. ",
        "conclusions": "We have shown that the recently discovered transient in M\\,82 is relatively faint in X-rays given its radio luminosity. The lack of variability in its radio light curve also make it unlikely that the source is a normal BH transient. The source could very well be the extragalactic analogue of SS\\,433 or a younger version of S\\,26. \\\\   \\textbf"
    },
    "1107/1107.3616_arXiv.txt": {
        "abstract": "We investigate the physical properties of dense cores formed in turbulent, magnetized, parsec-scale clumps of molecular clouds, using three-dimensional numerical simulations that include protostellar outflow feedback. The dense cores are identified in the simulated density data cube through a clumpfind algorithm. We find that the core velocity dispersion does not show any clear dependence on the core size, in contrast to Larson's linewidth-size relation, but consistent with recent observations. In the absence of a magnetic field, the majority of the cores have supersonic velocity dispersions. A moderately-strong magnetic field reduces the dispersion to a subsonic or at most transonic value typically. Most of the cores are out of virial equilibrium, with the external pressure dominating the self-gravity. The implication is that the core evolution is largely controlled by the outflow-driven turbulence. Even an initially-weak magnetic field can retard star formation significantly, because the field is amplified by the outflow-driven turbulence to an equipartition strength, with the distorted field component dominating the uniform one. In contrast, for a moderately-strong field, the uniform component remains dominant. Such a difference in the magnetic structure is evident in our simulated polarization maps of dust thermal emission; it provides a handle on the field strength. Recent polarization measurements show that the field lines in cluster-forming clumps are spatially well-ordered. It is indicative of a moderately-strong, dynamically important, field which, in combination with outflow feedback, can keep the rate of star formation in embedded clusters at the observationally-inferred, relatively-slow rate of several percent per free-fall time. ",
        "introduction": "\\label{sec:intro}  Millimeter and submillimeter observations of dense cores in nearby parsec-scale cluster-forming clumps have shown that the core mass function (CMF) resembles the stellar initial mass function (IMF). For example, \\citet{motte98} performed 1.2 mm continuum observations toward L1688 and identified 57 starless dense cores whose mass function ($dN/dM \\sim M^{-2.5}$) is consistent with the Salpeter power-law IMF at the high mass end ($dN/dM \\sim M^{-2.35}$). The CMF also has a break at about 0.3 M$_\\odot$, below which it flattens to $dN/dM \\sim M^{-1.5}$. Other authors have found similar CMFs in the region \\citep[e.g.,][]{johnstone00,stanke06,maruta10}. The shape of the CMF is broadly consistent with the IMF of Class II YSOs in the same region \\citep{bontemps01}. Another example is the Serpens Cloud Core, where \\citet{testi98} identified, through 3 mm continuum interferometric observations, 32 dense cores whose mass function can be fitted by a power-law of $dN/dM \\sim M^{-2.1}$, again consistent with the Salpeter IMF. These and other observations suggest that the identified dense cores may be the direct progenitors of individual stars and the bulk of the stellar IMF may be at least partly determined by cloud fragmentation in the parsec-scale dense clumps. Thus, understanding the formation process of dense cores is a key step towards a full understanding of how stars form.     Both supersonic turbulence and magnetic fields are expected to play a role in clustered star formation. The relative importance of the two is still under debated, however. One school of thought is that the core (and thus star) formation in a cluster-forming clump is mainly regulated by supersonic turbulence. In this scenario, star formation is completed rapidly, before the initial supersonic turbulence has decayed significantly \\citep[e.g.,][]{elmegreen00,hartmann01,klessen00}. Magnetic fields are less important in this picture, although \\citet{padoan02} argued that a weak magnetic field is still needed in order to produce a core mass spectrum that resembles the stellar IMF. The weak magnetic field is expected to be strongly tangled by the supersonic turbulence. Ordered field structures, if present at all, are typically attributed to large-scale compression and are expected to be parallel to the dense filaments that are also produced by the same compression \\citep[e.g.,][]{pelkonen07}.  The second school of thought envisions a magnetic field that is more dynamically important; it prevents stars from forming too rapidly. Because supersonic turbulence dissipates rapidly, it must be somehow replenished. In a cluster-forming clump, protostellar outflow feedback can play an important role in turbulence regeneration \\citep[e.g.,][]{maury09,arce10,nakamura11a,nakamura11b}. \\citet{li06} demonstrated that protostellar outflows can resupply the supersonic turbulence, keeping the clumps near a quasi-virial equilibrium state for a relatively long time \\citep[see also][]{matzner07,nakamura07,carroll09}. Because the matter moves preferentially along the dynamically important magnetic field, the field is expected to be more or less perpendicular to the dense filaments that are created by either turbulence compression or self-gravity. A goal of this paper is to quantify the effects of the magnetic field on the clustered star formation in general, and the properties of dense cores in particular, through 3D MHD simulations that include outflow feedback for turbulence replenishment.   The rest of the paper is organized as follows. In Section \\ref{sec:model}, we describe the numerical method and simulation setup. In Section \\ref{sec:results}, we apply a version of the so-called ``clumpfind'' algorithm to identify dense cores in a parent dense clump and derive their properties. We find that the internal motions of the cores are sensitive to the magnetic field strength of the parent clump and that the external surface pressures play a dominant role in core formation in outflow-driven turbulence. In Section \\ref{sec:discussion}, we compare the properties of the simulated cores with those observed in a couple of nearby cluster forming regions. Our main conclusions are summarized in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}    \\subsection{Comparison with Core Observations} \\label{subsec:observations}   Here, we compare the core properties obtained in Section \\ref{sec:results} with observations. We use the core data toward two nearby parsec-scale cluster-forming clumps: the $\\rho$ Ophiuchus Main Cloud and NGC 1333 in Perseus. In the  $\\rho$ Ophiuchus region, \\citet{maruta10} identified 68 H$^{13}$CO$^+$ ($J=1-0$) cores by applying clumpfind to the 3D position-position-velocity cube data obtained using Nobeyama 45~m telescope. In NGC~1333, \\citet{walsh07} identified 93 N$_2$H$^+$ ($J=1-0$) dense cores, again using clumpfind. Since both the H$^{13}$CO$^+$ ($J=1-0$) and  N$_2$H$^+$ ($J=1-0$) transitions have critical densities of order $\\sim 10^5$ cm$^{-3}$, the H$^{13}$CO$^+$ and N$_2$H$^+$ cores should have densities comparable to our simulated cores, which were identified using a threshold density $\\sim 10^5$ cm$^{-3}$. Since the moderately-strong field model yields a core velocity dispersion that is in better general agreement with observations than the non-magnetic or weak-field model, we will concentrate on comparing its cores with the $\\rho$ Oph and NGC~1333 data.  In Fig. \\ref{fig:obs}a, we compare the 1D FWHM nonthermal velocity width-radius relations for the $\\rho$ Oph cores of \\citet{maruta10} and our simulated cores (see Section \\ref{sec:results}). The 3D velocity dispersions discussed in Section \\ref{sec:results} was converted into the 1D FWHM velocity widths using $dV_{\\rm NT} = (8\\ln2/3)^{1/2}\\Delta V_{\\rm NT}$. Our simulated cores tend to be somewhat smaller than the $\\rho$ Oph cores; the average radius of the $\\rho$ Oph cores is estimated at 0.04 pc, about 1.7 times that of model S1. Aside from the size difference, the observed and simulated data are almost indistinguishable. Neither shows a clear dependence of linewidth with size, except for a weak trend for the linewidth to increase slightly with radius. Neither follows the Larson's linewidth-size relation. The observations are also consistent with the cores formed from initially-magnetically subcritical clouds \\citep[see Fig. 18 of ][]{nakamura08}. We did not plot the NGC 1333 cores on Fig. \\ref{fig:obs}a because the core radius in Table 1 of \\citet{walsh07} was quoted in single significant digit, which artificially introduces a large scatter in the plot. Nevertheless, inspection by eye reveals no clear correlation between the velocity width and core radius, broadly consistent with the data for both the $\\rho$ Oph cores and the simulated cores of model S1.    Figure \\ref{fig:obs}b compares the virial parameter-mass relations of the simulated and observed cores. The red and blue dots, and black crosses denotes the $\\rho$ Oph \\citep{maruta10} and NGC 1333 cores \\citep{walsh07}, and the cores of model S1, respectively. For the NGC 1333 cores, we recalculated the virial masses by including the thermal contribution from the gas with a temperature of 15 K and adopting a distance of 235 pc \\citep{hirota08}, instead of 300 pc used in \\citet{walsh07}. We also recalculated the LTE masses by adopting the N$_2$H$^+$ fractional abundance relative to H$_2$ of $X_{\\rm N_2H^+} \\sim 3\\times 10^{-10}$ derived from observations toward $\\rho$ Oph because \\citet{walsh07} determined their adopted N$_2$H$^+$ fractional abundance such that the points are distributed evenly about the line of equality ($M_{\\rm LTE}=M_{\\rm vir}$) in their $M_{\\rm LTE}$-$M_{\\rm vir}$ plot. This N$_2$H$^+$ fractional abundance is determined by taking average of the values obtained toward Oph A \\citep{james04} and B2 subclumps \\citep{friesen10}. Figure \\ref{fig:obs}b indicates that the virial parameter-core mass relation of the nearby cluster-forming clumps follows roughly a single power-law of $\\alpha_{\\rm vir} \\propto M^{-2/3}_{\\rm c}$ with a small dispersion, which is again in good agreement with model S1 that has a moderately-strong magnetic field. A dynamically important magnetic field may have played a role in core formation in these regions.    Besides the $\\rho$ Oph region, \\citet{maruta10} found that the velocity dispersions of H$^{13}$CO$^+$ cores are almost independent of the core radius in Orion A as well, even though the latter is more than a factor of 3 farther away. The coarser spatial resolution for Orion means that its cores may contain (smaller, $\\rho$ Oph-like) cores blending together. If this is true, then the nearly flat velocity-width-radius relation observed in both regions suggests that the inter-core motions among the neighboring cores are almost comparable to the internal motions in the individual cores. Such a feature was pointed out by \\citet{andre07} who measured the velocity difference among neighboring cores using N$_2$H$^+$ ($J=1-0$). If the velocity-width-radius relation is flat ($\\Delta V \\propto$ constant) and the core mass is proportional to $R_{\\rm c}^3$, then the virial parameter is scaled as $\\alpha_{\\rm vir} \\propto R_{\\rm c} \\Delta V^2 / (GM_{\\rm c}) \\propto M_{\\rm c}^{-2/3}$, a similar power-law to that derived by \\citet{bertoldi92}. This again suggests the importance of ambient turbulent pressure in dynamics of the cores as discussed in Section \\ref{sec:results}.     \\subsection{Amplification of Magnetic Field} \\label{subsec:amplification}   To isolate the effect of magnetic field on clustered star formation, we present in Fig. \\ref{fig:sfe}a the time evolution of star formation efficiency (SFE) for three models with different field strengths but no outflow feedback. Our simulations show that, as expected, the magnetic field tends to slow down star formation, especially at relatively late times. The rate of star formation remains rather high, however, particularly at early times. For example, from the formation of the first star to the time when the SFE reaches $\\sim 16 \\%$, the average star formation efficiency per global free-fall time (SFR$_{\\rm ff}$; Krumholz \\& McKee 2005) is estimated  to be 29 \\%, 27 \\%, and 21 \\%, respectively, for models N2, W2, and S2. These values are larger than the observationally inferred values of a few \\% \\citep{krumholz07}. The large SFR$_{\\rm ff}$ in the absence of outflow feedback is in agreement with previous studies (see also the Appendix for the analytic formula of SFR$_{\\rm ff}$). For example, using MHD SPH simulations, \\citet{price08} and \\citet{price09} followed the evolution of 50 $M_\\odot$ clumps with $0.25 \\le \\beta \\le  \\infty$ until $t\\lesssim 1.5 t_{\\rm ff,cl}$. Although their initial model clump mass is too small compared to nearby cluster-forming clumps like $\\rho$ Oph and NGC 1333 (which have masses of order $10^3$ M$_\\odot$), their star formation efficiencies per free-fall time are estimated to be over 10 \\%.  These simulations imply that it is difficult to reduce the star formation rate to the observed level by a moderately-strong magnetic field alone, and other factors are needed to significantly retard star formation in cluster-forming clumps.  One way to slow down star formation is the feedback from forming stars. \\citet{price09} found that radiative feedback from forming stars does not change the global star formation efficiency in a parent clump much, although it significantly suppresses the small-scale fragmentation by increasing the temperature in the high-density material near the protostars. Figure \\ref{fig:sfe}b shows that the star formation rate is greatly reduced by the inclusion of outflow feedback, even in the presence of a relatively weak magnetic field, in agreement with previous studies \\citep{li06,nakamura07,wang10}. The average star formation efficiency per global free-fall time is estimated to be 8 \\%, 6 \\%, and 4 \\% for model N1 (no magnetic field), W1 (weak field), and S1 (moderately-strong field), respectively. For the model with the moderately-strong field, the SFR$_{\\rm ff}$ is more or less comparable to the observed value.   The reason for the large reduction in SFR$_{\\rm ff}$ can be seen in Fig. \\ref{fig:magnetic field}, where the magnetic energy is plotted against the evolution time, for models W1 and S1. As shown in Fig. \\ref{fig:magnetic field}a, for the moderately-strong field case, the total magnetic energy is dominated by the background uniform field that does not contribute to the force balance in the initial cloud.  We therefore illustrate in Fig. \\ref{fig:magnetic field}b the time evolution of the magnetic energy stored in the distorted component that was amplified by supersonic turbulence. Here, we computed the magnetic energy stored in the distorted component by subtracting the initial magnetic energy from the total magnetic energy. For comparison, we also plotted in Fig. \\ref{fig:magnetic field}b the evolution of the total kinetic energy of the clump. Since the kinetic energy is often dominated by the unbound high-velocity gas associated with the active outflows, we approximate the total kinetic energy of the clump as the kinetic energy of the gas whose velocity is smaller than 10 $c_s$. For both models, the amplified components tend to increase with time. For the model with the moderately-strong field, the amplified component begins to oscillate around a level value after a free-fall time. For both models, the magnetic energy of the amplified component becomes comparable to the kinetic energy of the dense gas by the end of the computation, indicating that a quasi-equipartition has been reached \\citep[see also ][]{federrath11}. For the weak field case, the amplified component is more important than the initial uniform field, resulting in a significantly-distorted magnetic field structure. Such a highly-distorted magnetic field can be seen in the 3D bird's eye view of the  density and magnetic field distributions of model W1, which is presented in the left panel of Fig. \\ref{fig:3d}. In contrast, the global magnetic field is well-ordered for the case of moderately-strong initial magnetic field, which is shown in the right panel of Fig. \\ref{fig:3d}.   \\subsection{Polarization Maps} \\label{subsec:polarization}   Polarization maps of submillimeter thermal dust emission have recently been obtained for nearby star forming regions \\citep[e.g.,][]{houde04,girart09,matthews09}. Here, we present the polarization maps derived from the simulation data for the two magnetically-supercritical models with different initial magnetic field strengths (models W1 and S1 which have the magnetic flux-to-mass ratios of $\\bar{\\lambda}=4.3$ and 1.4, respectively). We computed the polarized thermal dust emission from the MHD model following \\citet{padoan01}. We neglect the effect of self-absorption and scattering because we are interested in the thermal dust emission at submillimeter wavelengths.  We further assume that the grain properties are constant and the temperature is uniform. The polarization degree is set such that the maximum is equal to 15 \\%. Figure \\ref{fig:polarization} shows the dust polarization maps calculated from models W1 and S1.    Only a small portion of the computation box is shown in each panel of Fig. \\ref{fig:polarization}. As expected from Fig. \\ref{fig:3d}, in the presence of a weak magnetic field, the spatial distribution of the polarization vectors has relatively large fluctuations. The polarization degree tends to be smaller in several parts where the magnetic fields are strongly distorted. In the densest parts, the polarization vectors appear more or less parallel to the local elongated structures or dense filaments, whereas there is no clear correlation between the vectors and density distribution in less dense parts. This may be due to the fact that the local dense filaments are created by turbulent compression that preferentially enhances the magnetic field component transverse to the shock plane. In contrast, in the presence of the strong magnetic field, the filamentary structure is prominent and the filament axes tend to be perpendicular to the polarization vectors that are almost parallel to the initial magnetic field direction, suggesting that the gas flow is channeled preferentially by the strong magnetic field to form the filaments. The polarization observations can thus provide a handle on the magnetic field strength of cluster-forming clumps.    Recent polarization observations of cluster-forming clumps show that the global magnetic field lines are more or less spatially well-ordered \\citep[e.g.,][]{davis00,houde04,sugitani10,sugitani11}. For example, \\citet{sugitani10} found that the Serpens cloud core is threaded by a hour-glass shaped well-ordered magnetic field and is elongated in the cross-field direction. \\citet{sugitani11} found that the Serpens South filamentary infrared dark cloud appears to be threaded by more or less straight global magnetic field. The observed spatially-ordered magnetic field structures imply that the magnetic fields in the nearby cluster-forming clumps are likely to be at least moderately-strong and dynamically important.   \\subsection{Core Mass Function} \\label{subsec:CMF}   In Figs. \\ref{fig:CMF}a, \\ref{fig:CMF}b, and \\ref{fig:CMF}c, we plot the core mass functions (CMFs) for our identified cores for models N1, W1, and S1, respectively. To enlarge the number of sample cores, we added up all the cores identified at the three stages when SFE has reached 0.08, 0.12, and 0.16. In addition, we added up the cores identified from the runs using the different initial turbulent realization for which the different random numbers were used for both the amplitude and phase. The total number of cores so identified is over 500 for each model.   The overall shapes of the CMFs are similar in all three cases, suggesting the shape of the CMF is relatively insensitive to the initial magnetic field strength. This is different from the turbulent fragmentation scenario proposed by \\citet{padoan01} where a weak magnetic field is required to reproduce the CMF that is similar in shape to the Salpeter IMF \\citep[see also][]{li10}. In our simulations, a moderately-strong magnetic field can produce a CMF that resembles the Salpeter IMF as well.  We note two interesting features of the computed CMFs. First, there appear to be a lack of massive cores compared to the Salpeter stellar IMF. This is not necessarily a problem, because it is unclear whether the nearby cluster-forming clumps such as $\\rho$ Oph and NGC 1333 that we aim to simulate would ever produce massive stars. Alternatively, massive stars can grow from initially less massive cores, fed by global collapsing flows towards the bottom of the clump potential (Smith et al. 2008; Wang et al. 2010). Second, the turnover at the low mass end is less clear in the CMFs than in the Kroupa or Chabrier IMF. This may be because, at the low mass end, most of the cores are gravitationally unbound, and may not form stars. Further studies are required to determine the relation between the CMF and IMF.      We have performed a set of 3D MHD simulations of cluster formation taking into account the effects of protostellar outflow feedback, and identified dense cores by applying a clumpfind algorithm to the simulated 3D density data cubes. The main results are as follows.   1. Dense cores do not follow Larson's linewidth-size relation. We find that the velocity dispersions of dense cores show little correlation with core radius, irrespective of the strength of the magnetic field and outflow feedback. In the absence of a magnetic field, the majority of the cores have supersonic velocity dispersions, whereas in the presence of a moderately-strong magnetic field, the cores tend to be subsonic or at most transonic.  2. We find that most of the cores are out of virial equilibrium, with the external pressure due to ambient turbulence dominating the self-gravity. The core formation and evolution is largely controlled by the dynamical compression due to outflow-driven turbulence. Such a situation is contrast to the strongly-magnetized (magnetically subcritical) case, where the self-gravity plays a more important role in the core dynamics, particularly for massive cores.   3. Even an initially-weak magnetic field can retard star formation significantly, because the field is amplified by supersonic turbulence to an equipartition strength. In such an initially weak field, the distorted field component dominates the uniform one. In contrast, for a moderately-strong field, the uniform component remains dominant. Such a difference in the magnetic structure can be observed in simulated polarization maps of dust thermal emission. Recent polarization measurements show that the field lines in nearby cluster-forming clumps are spatially well-ordered, indicative of a moderately-strong, dynamically-important, field."
    },
    "1107/1107.3091_arXiv.txt": {
        "abstract": "The efficiency of the Kolmogorov-Arnold technique for the astrophysical signals is studied modeling sequences with both random and regular properties. This technique has been applied to the study of the structures in cosmic microwave background maps obtained by the Wilkinson Microwave Anisotropy Probe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.0601_arXiv.txt": {
        "abstract": "{Measures of the \\hi\\ properties of a galaxy are among the most sensitive interaction diagnostic at our disposal. We report here on a study of \\hi\\ profile asymmetries (e.g., lopsidedness) in a sample of some of the most isolated galaxies in the local Universe. This presents us with an excellent opportunity to quantify the range of intrinsic \\hi\\ asymmetries in galaxies (i.e., those not induced by the environment) and provides us with a zero-point calibration for evaluating these measurements in less isolated samples.} { We aim to characterize the \\hi\\ profile asymmetries in a sample of isolated galaxies  and  search for correlations between \\hi\\ asymmetry and their environments, as well as their optical and far infrared (FIR) properties.} { We use high signal-to-noise global \\hi\\ profiles for galaxies in the AMIGA project (Analysis of the Interstellar Medium of Isolated GAlaxies, {\\tt http://amiga.iaa.es}). We restrict our study to  $N$ = 166 galaxies (out of 312) with  accurate measures of the \\hi\\ shape properties. We quantify asymmetries using a flux ratio parameter. } { The asymmetry parameter distribution of our isolated sample is well described by a Gaussian model. The width of the distribution is $\\sigma$ = 0.13, and could be even smaller ($\\sigma$ = 0.11) if instrumental errors are reduced. Only 2\\% of our carefully vetted isolated galaxies sample show an asymmetry in excess of 3$\\sigma$. By using this sample we minimize environmental effects as confirmed by the lack of correlation between \\hi\\ asymmetry and tidal force (one-on-one interactions) and neighbor galaxy number density. On the other hand, field galaxy samples show wider distributions and deviate from a Gaussian curve. As a result we find higher asymmetry rates ($\\sim$ 10--20\\%) in such samples. We find evidence that the spiral arm strength is inversely correlated with the HI asymmetry. We also find an excess of FIR luminous galaxies with larger HI asymmetries that may be spirals associated with hidden accretion events. } { Our sample presents the  smallest fraction of asymmetric \\hi\\ profiles compared with any other yet studied. The width of the associated asymmetry parameter distribution can help to distinguish the frequency and processes of self-induced HI asymmetries, and serve as a baseline for studying asymmetry rates in other environments.} ",
        "introduction": "\\label{sec:introduction}   The geometry and kinematics of gaseous disks in galaxies are mainly governed by the  gravitational potential of stellar and non-baryonic components. Any perturbation of the equilibrium caused by either internal or external processes, may produce asymmetries in these disks. Atomic gas (\\hi ) is the most extended cold component of the interstellar medium (ISM) and is a sensitive diagnostic of perturbations. The \\hi\\ component in spiral galaxies has long been known to show both geometrical  and kinematic asymmetries \\citep[e.g.,][]{1969Natur.221..531B,1972A&A....17..207H,1973A&A....29..447A,1980MNRAS.193..313B}. Asymmetries in stellar disks are also common and are traced by optical and near-infrared light (the latter less affected by dust extinction). Observations show that 30\\% of galaxies are significantly \\emph{lopsided} at near-infrared wavelengths \\citep{1994A&A...288..365B,1995ApJ...447...82R,1997ApJ...477..118Z,2005A&A...438..507B}. Asymmetries in the stellar component are not necessarily correlated with lopsidedness in the gaseous component \\citep{2000AJ....120..139K,2004AJ....127.1900W}. This lack of correlation is not surprising because the neutral hydrogen component in a galaxy is typically twice as extended as the stellar component \\citep[e.g.,][]{1994AAS..107..129B}, and might be perturbed in different ways with respect to the stellar component (e.g., stripping). The study of \\hi\\ properties is thus a better probe of past and/or recent perturbations than the stellar counterpart, especially for weak interactions.  Study of global \\hi\\ velocity profiles of galaxies has proven to be very useful for a quantification of the frequency and amplitude of disk asymmetries \\citep[e.g.][]{1994A&A...290L...9R}. While only aperture synthesis can provide full 2D information about the \\hi\\  distribution and kinematics, the 1D profiles provide a valuable measurement at a small fraction of the cost in observing time.   Past work suggests that the \\hi\\  asymmetry properties of galaxies do not depend strongly on local environmental conditions. Studies of  \\emph{field} and/or isolated galaxies suggest that \\emph{at least 50~\\%} show significant \\hi\\ profile asymmetries \\citep{1994A&A...290L...9R,1998AJ....115...62H}, and even higher $\\sim$75\\% in late-type spiral galaxies \\citep{1998AJ....116.1169M}. Although homogeneous studies of the asymmetry rate in richer environments are rare, it is usually believed that they show a comparable rate of asymmetric \\hi\\ profiles.   The implications of a high asymmetry rate essentially independent of environment is currently the subject of debate \\citep[for a review see][]{2009PhR...471...75J,2008A&ARv..15..189S}. It has been  suggested that the mechanism responsible for producing asymmetric disks must be long-lived because high asymmetry rates are observed in samples of field and/or isolated galaxies. Because the signatures of tidal encounters are relatively short-lived, lasting only on the order of a dynamical time-scale for a wide range in mass ratios, orientations, inclinations, relative velocities and impact parameters  (e.g. \\citealt{2005A&A...438..507B}), it cannot be the only agent responsible for the high asymmetry rate in different environments.  A number of longer-lived  mechanisms have been proposed, a) intermittent minor mergers \\citep{1996ApJ...460..121W,1997ApJ...477..118Z}, b) high-velocity cloud/gas accretion \\citep{2005A&A...438..507B,2008A&ARv..15..189S,2008arXiv0810.5130M}, c) halo-disk misalignment  \\citep{1998ApJ...496L..13L,2001MNRAS.328.1064N} and/or d) internal perturbations including sustained long-lived lopsidedness owing to non-circular motions \\citep{1980MNRAS.193..313B} or global m=1 instabilities \\citep{2007MNRAS.382..419S}.  In order to address the relevance of the different proposed mechanisms (internal versus external, short versus long-lived) one must first study a sample of well isolated galaxies in the nearby Universe ($\\lesssim$ 150 Mpc). This approach should minimize any contribution from tidal interactions and facilitate the interpretation of results with respect to other samples of galaxies. Reference samples used to study the rate of \\hi\\ asymmetries involve galaxies which, although assumed to be field/isolated, usually include a significant population of interacting galaxies \\citep[e.g.][]{1994A&A...290L...9R}. Note that field galaxies are defined as galaxies not belonging to the cluster environment and a significant number of them are  likely to be members of interacting pairs or multiplets  \\citep{2006A&A...449..937S}. \\citet{1994A&A...290L...9R} find that about half of the nearby field galaxies show asymmetric profiles, estimated from a compilation of six \\hi\\ surveys (1371 profiles were classified). It is important to remember that, 1) the overall environmental properties of the samples were not assessed meaning that  a significant fraction of galaxies might be environmentally influenced, and 2) asymmetries were assessed using qualitative criteria  \\citep{1994A&A...290L...9R}.  Statistical studies  of \\hi\\ asymmetries in large samples of galaxies selected according to a well defined isolation criterion and using an objective quantification are rare. The only existing systematic study of \\hi\\ asymmetries in an isolated  sample is that of  \\citet{1998AJ....115...62H}, who studied the asymmetry rate for $N$ = 104 ($N$ = 78) galaxies that obey a 0.5$\\arcdeg$ (1$\\arcdeg$ ) projected separation criterion with respect to any known companion in the Arecibo General Catalog  (AGC, private database of R. Giovanelli and M. P. Haynes), where 0.5$\\arcdeg$ corresponds to 175~kpc at a typical velocity of the core sample at V$_r$= 1500~\\kms . About half of the galaxies appear to show significant \\hi\\ asymmetries. However, only companions with a velocity difference $\\Delta$V$<$ 400~\\kms\\ relative to the primary are considered. The AGC is complete up to $m$~$\\sim$~15.4~mag and/or a diameter of 1\\arcmin\\ (equivalent to a typical linear size of  6~kpc). The velocity criterion could make this sample biased against unbound plunging encounters, and owing to the size limit of the catalog, small galaxies would not have been taken into account -- hierarchical systems of small galaxies may in principle produce significant asymmetries.  Other studied samples are composed of field galaxies. \\citet{1998AJ....116.1169M} studied a sample of $N$ = 30 moderate to low surface brightness late-type spirals. Their galaxies lie between 2$\\arcdeg$ and 6$\\arcdeg$ from the center of the Fornax cluster or are field galaxies.  They suggest that 77\\% of the \\hi\\ profiles in their sample show a relevant asymmetry. \\citet{2005A&A...438..507B} studied $N$ = 76 galaxies based on the OSUBS Galaxy Survey \\citep{2002ApJS..143...73E}, a sample not selected according to any environmental criterion, which hence may have a relevant amount of interacting galaxies. However, it has been recently referred to as a field galaxy sample \\citep[e.g.,][]{2009PhR...471...75J}. The asymmetry is larger than 10\\% for nearly 66\\% of the galaxies in this sample.   In this paper we use a large and complete sample of isolated galaxies \\citep{2005A&A...436..443V,2007A&A...462..507L} to evaluate the intrinsic distribution of \\hi\\  asymmetries. Comparisons show that our sample is more isolated than those used in previous studies and indeed is representative of the most isolated galaxies in the local Universe. The AMIGA project (Analysis of the interstellar Medium of Isolated GAlaxies\\footnote{\\tt http://amiga.iaa.es}, \\citealt{2005A&A...436..443V}) involves vetting and analyzing the properties of galaxies in the Catalogue of Isolated Galaxies ($N$ = 1050 galaxies, CIG, \\citealt{1973AISAO...8....3K}), and provides a good starting point for this aim. This project includes a refinement of the sample through 1) revision of optical positions \\citep{2003A&A...411..391L}, 2) analysis of optical properties and completeness \\citep{2005A&A...436..443V}, 3) revised optical  morphologies \\citep{2006A&A...449..937S} and 4) reevaluation of the isolation degree \\citep{2007A&A...470..505V,2007A&A...472..121V}. We use all these refinements in the present paper. The isolation criterion used for this compilation minimizes the probability of a major interaction within the last  $\\sim$ 3~Gyr \\citep{2005A&A...436..443V} while quantifying possible minor interactions. A multiwavelength characterization of different interstellar medium (ISM) components/phases and of the stellar component has been carried out including a) optical  \\citep{2005A&A...436..443V}, b) FIR \\citep{2007A&A...462..507L}, c) radio-continuum \\citep{2008A&A...485..475L}, and d) H$\\alpha$ emission \\citep{2007A&A...474...43V}, as well as e) nuclear activity \\citep{2008A&A...486...73S,2011A&A...0S}.  This paper presents an analysis of global (1D) asymmetry measures for $N$ = 312 \\hi\\ AMIGA galaxies with high signal-to-noise (S/N)   spectra (Sect.~\\ref{sec:sample}). We cleaned the sample of sources with uncertain asymmetry measures, yielding a total of $N$ = 166 galaxies with high reliability data. For this subsample we studied the \\hi\\ profile asymmetry rate (Sect.~\\ref{sec:lopsidedness}), the role of the environment for the rate of lopsided profiles (Sect.~\\ref{sec:environment}), as well as the correlations between \\hi\\ asymmetry and stellar properties (optical luminosity, morphological type and signs of perturbation), and star formation rate (as traced by FIR luminostity, Sect.~\\ref{sec:optical}). Finally we identify the underlying distribution of intrinsic profile asymmetries and discuss the possible origin of asymmetries in gaseous disks (Sect.~\\ref{sec:discussion}). ",
        "conclusions": "We used \\hi\\ global velocity profiles for a large sample of  isolated galaxies to $i)$ quantify the rate and amplitude of \\hi\\ asymmetries in isolated spiral galaxies, where environmental processes such as tidal interactions and ram pressure are minimized, $ii)$ study the role of the environment on the \\hi\\ asymmetries, and $iii)$ study possible correlations between \\hi\\ lopsidedness and the properties of the stellar component, including their morphological types, signatures of optical perturbation, bar and spiral strengths, as well as optical and FIR luminosities.  To quantify the \\hi\\ asymmetry, we calculated a flux ratio asymmetry parameter ($A_{flux~ratio}$). We restricted our study to a sample of $N$ = 166 galaxies (the \\hi\\ refined subsample) for which we minimized undesired artificially induced lopsidedness by avoiding large uncertainties owing to the rms of the profile, determination of the mean velocity, and pointing offsets.  We found that a half-Gaussian curve properly fits the $A_{flux~ratio}$ distribution of this refined sample, with a $\\sigma$ = 0.13. We suggest that if we deconvolve other sources of errors such as baseline fitting and random pointing offsets, then the underlying $\\sigma$ is reduced  to $\\sigma$ $\\sim$ 0.11. We confirm that by using this sample we effectively minimize nurture effects, because there is a lack of correlation between  \\hi\\ asymmetries and isolation parameters such as tidal force (one-on-one interactions) and number density.  We compared the distribution of \\hi\\ asymmetries of previously studied field galaxies with that of our isolated galaxy sample. A half-Gaussian fit does not successfully reproduce in general the asymmetry distribution of field samples. Indeed, the intrinsic $\\sigma$ is larger in field samples than in isolated galaxies. This is likely a result of the lack of an isolation criterion in the selection of field galaxies, which are likely contaminated by interacting objects. This suggests that environmental mechanisms (producing short-lived effects $\\sim$ 1~Gyr) are fundamental mechanisms to produce \\hi\\ asymmetries in galaxies, and are indeed responsible for the $\\sim$ 10--20\\% difference in the \\hi\\ asymmetry rate we see in field galaxies with respect to isolated galaxies. The asymmetry distribution of galaxies in denser environments is likely even wider and more skewed.  Within the isolated galaxy sample, we did not find any strong correlation between the \\hi\\ asymmetry and internal properties such as the morphological type, optical and FIR luminosities or signature of interaction. A signature of perturbed optical emission is not a necessary condition for the \\hi\\  profile to be asymmetric, and vice versa. We found a trend for larger \\hi\\ asymmetries to be located in more FIR luminous galaxies that are likely interacting objects. We also found evidence that galaxies with higher spiral arm strength have lower \\hi\\ asymmetries.  The here presented \\hi\\ refined subsample can be used to study the origin of intrinsic \\hi\\ asymmetries in isolated galaxies, and it is also a baseline for samples of galaxies in denser environments with \\hi\\ data that are properly evaluated for instrumental effects. This can help to shed light into the relative importance of different environmental and internally generated processes in shaping the \\hi\\ disks of galaxies."
    },
    "1107/1107.2929_arXiv.txt": {
        "abstract": "An orbiting object in a gas rich environment creates a gravitational density wake containing information about the object and its orbit. \\oops{Using linear perturbation theory,} we analyze the observable properties of the gravitational wake due to the object circularly moving in a static homogeneous gaseous medium, in order to derive the Bondi accretion radius $r_B$, the orbital distance $r_p$, and the Mach number $\\mach$ of the object. Supersonic motion, producing a wake of spiral-onion shell structure, exhibits a single-armed Archimedes spiral \\oops{and two-centered circular arcs with respect to the line of sight}. The pitch angle, arm width, and spacing of the spiral pattern are entirely determined by the orbital distance $r_p$ and Mach number $\\mach$ of the object. The arm-interarm density contrast is proportional to $r_B$, decreasing as a function of distance with a power index of $-1$. \\oops{The background density distribution is globally changed from initially uniform to centrally concentrated.} The vertical \\oops{structure} of the wake is manifested as circular arcs with the center at the object location. The angular extent of the arcs is determined by the Mach number $\\mach$ of the object motion. \\oops{Diagnostic probes of nonlinear wakes such as} a detached bow shock, an absence of the definite inner arm boundary, the presence of turbulent low density eddies, and elongated shapes of arcs are \\oops{explained in the extension of the linear analysis}. The density enhancement at the center is always $r_B/r_p$ independent of the nonlinearity, suggesting that massive objects can substantially modify the background distribution. ",
        "introduction": "\\label{sec:int}  The gravitational interaction of orbiting objects with a background medium is fundamental in many astrophysical problems. Objects in gas-rich environments such as planets in protoplanetary disks, young stars in natal molecular clouds, substellar mass objects in stellar winds, compact objects orbiting around supermassive black holes, and supermassive black hole binaries in colliding galaxies are subject to the gravitational interaction with background gas. The orbiting object naturally generates a gravitational density wake in its environment, inevitably followed by the reaction of the induced wake as dynamical friction.  A recent series of studies on the gaseous dynamical friction is devoted to understand the orbital evolution of the object in rectilinear motion \\citep{ost99,san99,kim09,nam10,nam11} and in circular motion \\citep{san01, kim07,kim08,kwt10}. These theoretical and numerical works have covered the effects of binarity, nonlinearity, acceleration, and background stratification. Albeit all these factors give rise to appreciable effects, the density wake and the dynamical friction force are primarily characterized by the accretional and orbital properties of the object described by the linear perturbation analyses in \\citet{ost99} and \\citet[hereafter \\kk]{kim07}.  Given the importance of the gravitational drag, these linear analyses have been applied to various astrophysical aspects such as the orbital decay of planets around giant stars \\citep{sch08,vil09}, the primordial mass segregation in massive star forming regions \\citep{gar10,cha10}, the binary hardening in parent gaseous disk \\citep{bar11}, the deceleration of galaxies as the heating sources of intracluster medium against the cooling flows \\citep{elz04,kwt05,kwt07,con08}, and the coalescence of supermassive black holes in gas-rich mergers of galaxies \\citep{esc04, esc05,dot06,dot07,may07}. These linear analyses have been extended to the relativistic case \\citep{bar07} for extreme mass-ratio inspirals in the vicinity of supermassive black holes. In particular, \\citet{bar08} found an important difference in the inclination angle of the orbits between the dynamical friction and the radiation reaction possibly detectable by Laser Interferometer Space Antenna (LISA).  Also important are the morphological and kinematic features of the density wake per se, as they can be directly observed by means of X-rays in shocked regions, molecular line emissions, dust continuum emissions, and in dust scattered light. \\citet{fur02} raised the possibility that the nature of dark matter can be examined by detailed X-ray observations of wakes due to the motion of galaxies in the intracluster medium, exhibiting distinguishable pictures between the standard collisionless cold dark matter and the collisional (fluid) dark matter. \\citet{san09} furthermore proposed that the morphology of a galactic wake in X-ray emission differs in purely baryonic Modified Newtonian Dynamics (MOND) from the case of Newtonian gravity with dark matter. These authors only inspected the local attributes of the gravitational density wakes using \\citeauthor{ost99}'s \\citeyearpar{ost99} formulae for the cases of the perturbing object in rectilinear motion.  However, many observable objects have fairly short orbital periods compared to the dispersion timescale of the background gaseous matter, leading to global shaping of the induced wakes. \\kk\\ \\oops{showed that} the curvilinear orbit modifie\\oops{s} not only the curvature of the wake but transforme\\oops{s} the overall structure entirely because of the memory effect\\footnote{\\ Memory effect (or battery effect) originally refers, only in nickel cadmium rechargeable batteries, to the symptom of losing their maximum energy capacity after being repeatedly recharged to partially discharged states. We apply this term here for the wake repetitively perturbed by the revolving object over the former perturbations remained, resulting in a loss of density contrast, as we will mention at the end of this paper.}. One application of the large scale wakes is to a case that planets or brown dwarfs lie in cool giant envelopes, whose wakes are possibly observable at large distances from their bright parent stars. Future high sensitivity/resolution observations of these wakes can provide evidence of the unseen substellar mass objects, further allowing us to derive their orbital properties (\\citealp{kim11} in prep.; for a brief introduction see \\citealp{kim10}).  Although \\kk\\ calculated the density wake, they primarily focused on the dynamical friction force that affects the initially circular orbit \\oops{and the description of the general wake feature was somewhat superficial. A more detailed understanding of the wake characteristics is necessary in order to provide an interpretative framework for both further theoretical works under more complicated situations and future observations to detect the gravitational wakes of hidden objects}. In this paper, we extend the semianalytic work \\oops{of \\kk}\\ highlighting the observable quantities, and determine the detailed structure and kinematics of the gravitational wake due to an object in circular motion. \\oops{Particularly, we study the density structure in great detail based on an analytic linear analysis, which also helps to understand aspects of the simulated density wake in a nonlinear regime. We provide the density contrast of the wake as a function of the object mass, orbital radius, orbital speed, and the background sound speed, providing a way to obtain the information of unseen orbiting object from the observation of the wake at a large distance.}  This paper is organized as follows. In \\S\\ref{sec:rev} we review the derivation and findings in previous works as the starting point of this investigation. In \\S\\ref{sec:lin} we present the nature of linear wake as revealed in the \\oops{shape} (\\S\\ref{sec:mor}), the density enhancement (\\S\\ref{sec:den}), and the velocity distribution (\\S\\ref{sec:vel}), provided that the weak gravitational influence is thoroughly described in the linear perturbation analysis. In \\S\\ref{sec:non}, based on the linear wake properties, we suggest a new explanation for nonlinear effects manifested on the wake when the perturbation in discontinuity is saturated. In \\S\\ref{sec:sum} we summarize the newly found properties of the gravitationally induced wake of a circularly orbiting object in a static uniform gaseous medium, and briefly consider applications to a few astronomical objects. ",
        "conclusions": "\\label{sec:sum}  The formation of a gravitational wake is a characteristic feature of an astronomical object moving through a medium. We examine the wake properties in detail and provide the observable aspects in a quantitative way, by revisiting the previous semianalytic and numerical works (\\kk; \\citealp{kwt10}) and comparing the features with the linear trajectory counterparts discovered in analytic and numerical methods \\citep{ost99,kim09}.  \\oops{T}he analytic formulae for the density wake in the previous work (reviewed in eq.~[\\ref{equ:k07}] and [\\ref{equ:sol}]) do not give a clear \\oops{physical} insight of spiral- and/or arc-shapes. Hence it is difficult to deduce the properties of the wake relevant to the object properties. A further calculation in this paper leads to \\oops{a} purely analytic solution for the wake shape (eq.~[\\ref{equ:mor}]), \\oops{which allows detailed description of the wake shape}. Simple geometrical approaches are adopted to provide the physical understanding of such shapes. Density enhancement profiles are also investigated. Since the empirical formulae based on the analytic analysis \\oops{(eq.~[\\ref{equ:alm}]--[\\ref{equ:alp}])} outline the maximum and minimum density enhancement very successfully, these can be directly adopted for interpretation of future simulation and observational data. Our deeper understanding of the linear wake helps to understand \\oops{nonlinear wakes shown in previous numerical simulations}. Particularly, nonlinear aspects such as the increase of background density, the disappearance of inner arm boundary, and the appearance of a detached spiral shock are predictable from a close inspection of linear wake properties.  \\oops{Considering} an object characterized by Bondi accretion radius $r_B$ \\oops{and} circularly orbiting at a distance $r_p$ with an orbital Mach number $\\mach$ in a static uniform gaseous environment\\oops{, o}ur new findings for linear wakes are as follows: \\begin{enumerate} \\item In the orbital plane of the perturbing object, the linear wake shows excess density in the form of single-armed structure. For supersonic motion, it is confined by distinct boundaries of Archimedes spirals in the form of $r/r_p=\\varphi\\mach^{-1}+C$ with different constants $C$'s for the outer (always 1) and the inner (depending on $\\mach$) boundaries. This expression for the explicit form of the spiral shape indicates the pitch angle of $\\tan^{-1}(\\mach\\,r/r_p)^{-1}$ and the arm spacing of $2\\pi r_p\\mach^{-1}$, from which we can deduce the orbital parameters, $\\mach$ and $r_p$. \\item The vertical structure of the gravitational density wake is quantified for the first time. In meridional planes, the dense regions are bounded by two-centered concentric arcs forming crescent-like structures ($\\mach>1$), partially overlapping each other for high orbital Mach number ($\\mach>4.6$) cases. The angular size of the arcs are up to $\\lesssim2\\tan^{-1}(\\mach^2-1)^{1/2}$. \\item The density enhancement $\\alpha$ at the orbital center is always $r_B/r_p$. \\item The density enhancement at the edges of the spiral arm pattern can be approximated to $\\alpha_{\\rm arm}=\\alpha_1+(4n+2)\\;\\amin$ with an integer number $n$ depending on $\\mach$ and some definitions of $\\alpha_1$ (eq.~[\\ref{equ:al1}]) and $\\amin$ (eq.~[\\ref{equ:alm}]). The interarm density also increases to $\\alpha_{\\rm inter}=(3n-2)\\;\\amin$ for supersonic cases ($n\\geq1$) and $\\amin$ for subsonic cases ($n=0$). Hence the density contrast between arm and interarm regions, $(\\rho_{\\rm arm}-\\rho_{\\rm inter})/\\rho_{\\rm inter}$ $\\approx$ $\\alpha_{\\rm arm}-\\alpha_{\\rm inter}$ in the linear regime ($\\alpha_{\\rm arm},\\;\\alpha_{\\rm inter}\\ll1$), is $\\sim\\alpha_1+(n+4)\\,\\amin$ for supersonic cases and $\\sim\\alpha_1+\\amin$ for subsonic cases. The numerical values for density contrast are approximately $r_B/r$ at large distance ($r\\gg r_p$), i.e., (0.7--2.9)\\,$r_B/r$ for a wide range of orbital Mach numbers $\\mach=0$--0.9 and 1.5--10. \\item The initially static fluid acquires velocity components in an approximate spherical form, starting from near the inner boundaries of high density regions and finally accumulating at the outer boundaries. \\item The gravitational size of the object does not significantly affect the linear wake features insofar as the size is smaller than the interarm gap. \\end{enumerate} As the accretion capability ($r_B/r_p$) increases for a marginally supersonic orbital Mach number ($\\mach$), the thermal pressure can be saturated at the junction of high density boundaries appearing as a tip in the linear analysis, so that the inner boundary extends beyond the outer boundary of the pattern. The induced wake reveals the following nonlinear behaviors: \\begin{enumerate} \\item An appearance of a detached bow shock and a disappearance of the definite inner arm boundary in the highly nonlinear wake are mutually explained by an overpass of the inner boundary beyond the outer boundary. The shock standoff distance depends only on $r_B$ and $\\mach$, regardless of the adiabatic index $\\gamma$. \\item Turbulent low density eddies are developed in the unstable process of producing the detached shock via the overpass of the inner boundary beyond the outer boundary. These eddies survive when the shock distance in the lateral direction is larger than $r_B\\mach^{-2}$, consistent to the linear trajectory counterpart in the nonlinear regime. \\item The density enhancement $\\alpha=r_B/r_p$ at the orbital center is preserved even in the extremely nonlinear regime, but the interarm density systematically increases proportional to this value, reducing the arm-interarm density contrast. \\end{enumerate}  Using these theoretical diagnostics, we can derive the object mass, velocity, and distance from the orbital center with a given sound speed of the medium. In the case that concentric arcs are observed, the perturbing object is expected to be at the center of the circular arcs, potentially giving $r_p$. The angular size of the arcs corresponding to $2\\tan^{-1}(\\mach^2-1)^{1/2}$ probes the orbital Mach number $\\mach$ of the object of interest, or the corresponding velocity $V_p$ with the aid of an estimate of the sound speed $\\cs$. The object mass $M_p$ (or the accretion radius $r_B=GM_p/\\cs^2$) can be roughly estimated from the density contrast $\\sim r_B/r$ at a distance $r$ from the center. High angular resolution observations are required to identify the broadening of arm edges ($\\sim2r_s$) so as to estimate the object size $r_s$. On the other hand, an observation of a spiral structure can provide information on $r_p$ and $\\mach$ by measurement of arm spacing and arm width. An observation of smooth density distribution without sharp discontinuity provides an upper limit of the object speed ($\\mach<1$). Finally, an observation of the arm structure with a missing inner boundary suggests the nonlinear wake condition, $r_B/r_p\\,(\\mach^2-1)^{-1}\\gtrsim0.1$.  From observations of spirals and/or arcs around stars, one may search for the possible existence of unseen companions, which are difficult to be directly detected either due to high obscuration by the dense environment or due to the bright glare of the central object. The presence of the companions may be inferred by their imprints on the environment at large distances from the star. For example, spiral structures have been observed around AB Aur revealed in near-infrared scattered emission \\citep{fuk04}, around AFGL 3068 in dust scattered galactic light \\citep{mau06}, and around CIT 6 in molecular line emission \\citep{din09}.  According to the result in this paper, the density contrast of a gravitational density wake due to a 10 Jupiter mass planet, \\oops{orbiting at a distance of several AU}, is expected to be of the order of 10\\,\\% at a distance of a few hundreds AU from the central star, where the spiral arms are located in the circumstellar disk of the Herbig Ae star AB Aur \\citep{fuk04}. Although the density contrast of the order of 10\\,\\% is not significant, it may be distinguishable in this disk (rather than spherical) system. Whether such structures can be observed for low mass objects remains to be ascertained. And whether the density contrast due to the gravitational influence of planets has a similar order of magnitude in a disk system also needs to be checked. To clarify the claim of \\citet{lin06} of forming giant planets as the origin of the observed spiral arms around AB Aur, it is necessary to carefully examine the density profile of the structure.  In the cases of AFGL 3068 and CIT 6, spiral structures are detected on the scale of 1--10 thousand AU \\citep{mau06,din09}, which exclude the possibility of planet wakes since the density contrast decreases with distance ($\\lesssim1$\\,\\%). The large spiral structures in these evolved circumstellar envelopes are likely to be caused by reflex motion of the mass losing star in comparable mass binary system \\citep{sok94,mas99,he07, edg08}. \\oops{To date, AFGL 3068 is the only object that shows a clear spiral modulation in the circumstellar envelope and its companion star is indeed confirmed by a direct imaging in near-infrared \\citep{mor06}. The binary separation is about a hundred AU.} Recently, \\citet{rag11} estimated approximately 1 solar mass for the companion of AFGL 3068 from the spiral shape. With this companion mass, the gravitational wake is expected to contribute as much as 100\\,\\% of the envelope density at \\oops{an observing} distance \\oops{of the wake} of a thousand AU, doubling the density in the arm structure. This density contrast of the companion wake may be comparable to the density enhancement due to the reflex motion of the primary star. Hence, in the central region ($\\lesssim$\\,1\\arcsec\\ by adopting the distance of 1\\,kpc to AFGL 3068), which will be soon resolved by the Atacama Large Millimeter/submillimeter Array (ALMA), there could exist two separated arms caused by two different mechanisms, i.e., one structure due to the reflex motion of the mass losing star and one gravitational density wake of the companion star. A comparison between the two mechanisms is needed in a range of parameter space relevant to the object and background properties.  We note that our analysis is limited to objects in circular motion in a static uniform isothermal gaseous medium. For comparison with detailed observations, our assumptions for the background medium and the object motion may need to be relaxed. However, the wake features described in this paper may suffice to provide a zeroth order approximation of the primary features expected from the gravitational wakes due to such objects."
    },
    "1107/1107.1211_arXiv.txt": {
        "abstract": "{We consider cosmological parameters estimation in the presence of a non-zero isocurvature contribution in the primordial perturbations. A previous analysis showed that even a tiny amount of isocurvature perturbation, if not accounted for, could  affect standard rulers calibration from  Cosmic Microwave Background observations such as those provided by the Planck mission, affect Baryon Acoustic Oscillations interpretation, and introduce biases in the recovered dark energy properties that are larger than forecasted  statistical errors from future surveys.  Extending  on this work,  here we adopt a general fiducial cosmology which includes a varying dark energy equation of state parameter and curvature. Beside Baryon Acoustic Oscillations measurements, we include the information from the shape of the galaxy power spectrum and  consider a joint  analysis of a Planck-like Cosmic Microwave Background probe and a future, space-based,  Large Scale Structure probe not too dissimilar from   recently proposed surveys.  We find that this allows one to break the degeneracies that affect the  Cosmic Microwave Background and Baryon Acoustic Oscillations  combination. As a result, most of the cosmological parameter systematic biases arising from an incorrect assumption on the isocurvature fraction parameter $f_{iso}$, become negligible with respect to the statistical errors.   We find that the Cosmic Microwave Background and Large Scale Structure combination gives a statistical error $\\sigma(f_{iso}) \\sim 0.008$, even when curvature and a varying dark energy  equation of state are included, which is smaller than the error obtained from Cosmic Microwave Background alone when flatness and cosmological constant are assumed. These results confirm the synergy and complementarity  between  Cosmic Microwave Background  and Large Scale Structure, and the great potential of   future  and planned galaxy surveys. }  \\begin{document} ",
        "introduction": "The standard cosmological model  assumes  adiabatic initial conditions for the perturbations. This simple adiabatic picture is well-motivated as it is predicted by the simplest inflationary  single-field models  \\cite{mukhanov-adiabatic,Brandenberger92}, and so far it provides an excellent fit to current data  (e.g., \\cite{KomatsuWMAP7}). There is however no a priori reason to discard more general initial conditions involving entropy isocurvature perturbations as they can arise from several different mechanisms e.g., multi-field inflation, \\cite{Linde1985,Kofman1986, Mollerach1990, Polarski1994, Langlois99, Peebles1999, Bartolo2001}, neutrino isocurvature perturbations \\cite{Bucher-general-2001}, axion dark matter \\cite{Axenides:1983hj, Lindeaxion, SeckelTurner, hybrid, TurnerWilczek, LindeLyth, Lyth:1991, Shellard:1997, Kawasaki1995axion} or the curvaton scenario \\cite{curvaton1, curvaton2, curvaton3}. Pure isocurvature models have been observationally excluded \\cite{Efstathiou86, Enqvist02, Page03, Hinshaw06},  but current observations still allow for an admixture of adiabatic and isocurvature contributions \\cite{KomatsuWMAP7, Beltran-04, Beltran-05,Valiviita09,Dunkley05,Keskitalo07}.  To date, Cosmic Microwave Background (CMB) data have been the main source of information not just about cosmological parameters but also about the nature of cosmological initial conditions. Relaxing the assumption of adiabaticity and allowing an isocurvature contribution introduces new degeneracies in the parameters space which weaken considerably  cosmological constraints e.g., \\cite{trotta2003, Valiviita09, Kurki05,Langlois00, Valiviita03, Bucher00,Sollom09}. Up to now,  the CMB temperature power spectrum alone has not been able to  break the degeneracy between the nature of initial  perturbations (i.e. the amount  and properties of an  isocurvature component) and cosmological parameters; adding external data sets somewhat alleviates this issue for some degeneracy directions e.g., \\cite{Beltran-04,Dunkley05}. As shown in \\cite{bucher-pol}, the precision polarization measurements of future and on-going CMB experiments like Planck will be crucial to lift such degeneracies.  Ref.~\\cite{Mangilli10} investigated the effect of relaxing the hypothesis of purely adiabatic initial conditions and found that even a tiny isocurvature fraction contribution, if not accounted for, can lead to an incorrect determination of the cosmological parameters and can affect the standard ruler calibration from the CMB. In fact,  the presence of an isocurvature component changes the shape and the location of the CMB acoustic peaks, mimicking the effect of  parameters such as $\\Omega_{m_0} h^2$, $H_0$ and $w$ (see  also \\cite{Valiviita09} and references therein). Ref.~\\cite{Mangilli10}  found that this has a crucial effect on ``standard rulers\", like the sound horizon at radiation drag, inferred from CMB observations. This is of relevance for the next generation of galaxy surveys which aims at probing with high accuracy the late time expansion and thus the nature of dark energy by means of Baryon Acoustic Oscillations (BAO) at low redshift ($z<2$). In fact, for Large Scale Structure (LSS), even a tiny isocurvature fraction contribution, if not accounted for,  can affect the  BAO estimation, introducing systematics in the BAO observables which can be larger than the statistical errors expected from future surveys.   BAO measurements from galaxy surveys have the potential to be an unprecedented powerful probe of the low redshift Universe, and thus of dark energy (\\cite{EisensteinHu, SE03, DETF} and references therein). The BAO in the primordial photon-baryon fluid, responsible for the characteristic peak structure of the CMB power spectrum, leave, in fact, an imprint in the large scale matter distribution and, at each redshift, their physical properties are related to the size of the sound horizon $r_s(z_d)$ at radiation drag (i.e., when baryons were released from the photons). Since the CMB in principle can precisely provide the standard ruler, i.e. $r_s(z_d)$, by measuring the BAO location at low redshift with LSS surveys, it is possible to probe the expansion history, i.e., the Hubble parameter $H(z)$, and the angular diameter distance $D_A(z)$ at different redshifts, and thus the dark energy properties.  It has been found that the effect of a number of theoretical systematics such as non-linearities, bias etc., on the determination of the BAO location can be minimized to the point that BAO are one of the key observables of the next generation dark energy experiments. However, it is important to investigate and test the robustness of the method to all possible theoretical uncertainties.  The analysis in Ref.~\\cite{Mangilli10} focused on the forecasts from a CMB-only Planck-like experiment to quantify the impact  of the isocurvature contribution on cosmological parameter estimation and degeneracies with $r_s(z_d)$ and, therefore, with BAO measurements. Here we extend the work to a joint analysis of a CMB Planck-like experiment combined to a future space-based LSS survey with characteristics not too dissimilar from the proposed Euclid mission\\footnote{http://sci.esa.int/euclid}. One of the key points is to see whether adding the full cosmological information provided by future LSS surveys could eventually break the degeneracies introduced by the presence of an isocurvature fraction in the initial conditions. In the present work in fact, we exploit, in the LSS case, not only the information enclosed in the  BAO positions, but include the galaxy power spectrum shape adopting the so-called ``$P(k)$-method marginalised over growth information'' \\cite{Melita11,Wang06}. Moreover we consider  spatial curvature and  a more general case of dark energy with a varying equation of state parametrized by $w_0$ and $w_a$.  The rest of the paper is organized as follows. In Sec.~2 we review theory and notation for isocurvature perturbations. In Sec.~3  we outline our methodology, paying particular attention to the method used to compute expected constraints from LSS surveys, including not only the BAO location measurements but also the power spectrum shape. We present our results in Sec.~4 and the conclusions in Sec.~5. ",
        "conclusions": "In this work we have relied on the Fisher matrix approach to forecasts, which is not well suited to quantify systematic effects. The errors we have reported are statistical only, which are optimistic in the sense that we assume that systematic errors are vastly sub-dominant.  Let us therefore discuss sources of systematic errors of major concern and their possible consequences on our results: these are the effects of non-linearities and galaxy bias.  In the  analysis presented here we have used the linear theory matter power spectrum and a scale-independent  galaxy bias. Our choice of the maximum k-mode ensures that the effects of non-linearity are quite mild and that eventually, before an analysis is performed on real data,  they could be  accurately modeled. Having considered only the mildly non-linear regime ensures that non-linearities will not erase or reduce significantly the signal. Reconstruction techniques    that  are being developed \\cite{reconstruction1,reconstruction2} and/or are already being applied to surveys \\cite{reid}  offer promising prospects. Thus while it will  be mandatory to include non-linearities in the actual data analysis, the forecasted errors are not made artificially smaller by using the linear matter power spectrum to compute our Fisher matrices.  Including smaller, more non-linear scales might help to improve the errors but would worsen the systematic effect of non-linearities and of a possible scale-dependent and/or non-linear bias. In this work we have made the simplifying assumption that bias is scale-independent but that the redshift dependence is not well known and so it is marginalised over. Bias is known to be scale-independent on large scales  but scale dependent for small scales and/or very massive halos. While halo bias can be modeled, the bias of galaxies hosted in dark matter halos can be more complicated especially at small scales. Here we have considered scales  where the so-called two halo term dominates and so complicated bias arising from the details of galaxy formation and   halo occupation distributions do not matter. In addition we have checked that most of the information, useful to break degeneracies between the sound horizon at radiation drag and  isocurvature contribution, come mostly from very large scales (those constraining the matter -radiation equality). The BAO signal is localized on scales where scale-dependent bias is more likely to be an issue, but the BAO location is relatively  robust to this.   Our reported error-bars could have been made smaller by considering also information about the growth of structures (i.e., redshift space distortions and the normalization of the power spectrum), however we decided to, conservatively, marginalise over the amplitude to be less sensitive to possible assumptions about galaxy bias.  As found in Ref.~\\cite{Mangilli10} for a Planck-like CMB experiment, even a tiny isocurvature fraction contribution, if not accounted for, can lead to an incorrect determination of the cosmological parameters even for a standard flat model with a cosmological constant. In this case however, when the isocurvature parameter $f_{iso}$ is added to the modelling, the errors increase but the cosmological parameters are recovered correctly. The error degradation from a CMB-only analysis is particularly severe for models where  spatial flatness is not imposed and curvature is left as a free parameter.  In this work we extend the analysis done in Ref.~\\cite{Mangilli10} and present the results from a joint analysis of a CMB Planck-like experiment and a LSS EULID-like galaxy survey, for a cosmological model that includes an isocurvature fraction in the initial conditions, varying dark energy and curvature.  As shown in Figs.~\\ref{fig_error2}-\\ref{fig_error} and in Tab.~\\ref{t:cmb+bao}, we find that adding the cosmological information coming from an Euclid-like LSS probe strongly reduces the degeneracies introduced by the presence of an isocurvature contribution with respect to a Planck-like CMB probe alone.  When isocurvature is included in the analysis, we find that the most affected parameter is the spectral index $n_s$ for which the degeneracy with the isocurvature fraction parameter $f_{iso}$ is never fully solved even when combining CMB and LSS. This is because in this paper we assume a curvaton underlying model as a working example for which $n_s \\equiv n_{ad}=n_{iso}$. Therefore $n_s$ encodes the information of both the spectral indices which are correlated with $f_{iso}$ for both CMB and LSS. In this case, from the combination CMB+LSS, we find an absolute $n_s$-shift$=0.277$ which is larger than the statistical 1--$\\sigma$ error by two order of magnitude. For all the cosmological parameters the correlations with $f_{iso}$ and the corresponding shifts are reported in Tab.~\\ref{t:cmb+bao}. Our findings are quantitative only for the  $n_{ad}=n_{iso}$ case, but can hold qualitatively for other models with $n_{iso} \\simeq 1$, as for example the axion-like isocurvature model. In this case, since $n_{ad}$ is close to scale invariant, we expect that the effect will be quantitatively very similar.  We investigate also the effect on the recovered cosmological parameters  in the case of the presence of a small amount of isocurvature  which is however ignored in the analysis. Ref.~\\cite{Mangilli10} found that even a tiny isocurvature fraction, if neglected, could bias the determination of ``standard rulers'' like the sound horizon $r_s(z_d)$ at radiation drag and introduce systematic biases  in the interpretation of BAO observables that are larger than the  statistical error-bars forecast for future and forthcoming surveys.  Ref.~\\cite{Mangilli10} considered a flat Universe where the dark energy is a cosmological constant, but in a more general cosmology the effect can be dramatically worst.  Here we find that adding LSS information to CMB priors eliminates the bias (i.e. reduces it well below the statistical errors) in the case of a small amount of isocurvature fraction $f_{iso} \\sim {\\cal O} (10^{-2})$, even for a general cosmological model with varying dark energy and curvature. The only parameter for which the systematic shift remains still comparable to the 1--$\\sigma$ error is the spectral index $n_s$.  These remarkable results are obtained  by better exploiting the potential cosmological information of  future LSS surveys by adding to BAO measurements also the power spectrum shape, following the so-called $P(k)$-method marginalised over growth information. In fact, the additional cosmological information in the power spectrum helps  to break degeneracies among the cosmological parameters, and thus indirectly improves the constraints on the isocurvature parameter $f_{iso}$.  Considering that -conservatively- we have not included information on the power spectrum amplitude and on the growth of cosmological structures (we marginalize over redshift space distortions parameters and LSS power spectrum normalization), we conclude highlighting once again the synergy and complementarity  between  CMB and LSS and the great potential of future galaxy surveys."
    },
    "1107/1107.4091_arXiv.txt": {
        "abstract": "We provide new measurements of wavenumbers and line identifications of \\nlines UI and UII near-infrared (NIR) emission lines between $2500$ cm$^{-1}$ and $12~000$ cm$^{-1}$ ($4000$ nm to $850$ nm) using archival FTS spectra from the National Solar Observatory (NSO). This line list includes isolated uranium lines in the Y, J, H, K, and L bands ($0.9$ $\\mu$m to $1.1$ $\\mu$m, $1.2$ $\\mu$m to $1.35$ $\\mu$m, $1.5$ $\\mu$m to $1.65$ $\\mu$m, $2.0$ $\\mu$m to $2.4$ $\\mu$m, and $3.0$ $\\mu$m to $4.0$ $\\mu$m, respectively), and provides six times as many calibration lines as thorium in the NIR spectral range.  The line lists we provide enable inexpensive, commercially-available uranium hollow-cathode lamps to be used for high-precision wavelength calibration of existing and future high-resolution NIR spectrographs. ",
        "introduction": "Astronomical spectrographs are used for a variety of high-precision measurements, ranging from the discovery of low mass exoplanets \\citep{2009A&A...507..487M} to the possible variation of fundamental constants, such as the fine structure constant \\citep{2001PhRvL..87i1301W} or the proton-electron mass ratio \\citep{MalecKeck2010MNRAS}. These works require excellent wavelength calibration sources and a detailed understanding of the associated uncertainties and systematics. In the era of extremely large telescopes (ELTs) it is often the accuracy of the calibration source, not the intrinsic photon noise, that limits the achievable precision. Furthermore, the science goals of future ELTs will require very high-precision calibration sources.  Below $900$ nm, the well-established thorium-argon (Th/Ar) hollow cathode lamps have been a workhorse~\\citep{1983ats..book.....P}. Continual improvements in the line list have now enabled Th lamps to be used to calibrate almost the entire optical bandpass with high precision \\citep{2007A&A...468.1115L, 2007MNRAS.380..839M}.  In contrast to the optical, the near infrared (NIR) does not yet have a widely-used calibrator like Th/Ar (see \\citet{2009ApJ...692.1590M} for a detailed discussion). \\citet{kerber2008th} have provided an atlas of Th/Ar lines in the NIR, which is sufficient for precisions of $\\approx 50$ m/s on CRIRES (CRyogenic high-resolution InfraRed echelle Spectrograph) \\citep{kaeufl2004crires}, but the density of Th lines is relatively low in the NIR, and many of the thorium lines are very weak.  Better precisions of $6$ m/s to $10$ m/s have been achieved by \\citet{2010A&A...511A..55F} with Telluric CO$_2$ lines, and $5$ m/s by \\citet{2010ApJ...713..410B} with an ammonia (NH$_\\textrm{3}$) gas cell in the K band.  Such calibration techniques are suitable for specific problems, but are only useful over specific ranges of the NIR.  An ideal emission calibration source would cover the observed spectrum in a forest of isolated lines of similar intensity, with the intrinsic wavelength of each line known to a high degree of precision. Laser frequency combs are such a calibration source~\\citep{2007MNRAS.380..839M, 2007SPIE.6693E..44O, 2008EPJD...48...57B, quinlan:063105}, but there are numerous technical hurdles before these systems will be considered turn-key: they are expensive to build and operate, they have yet to demonstrate the long-term stability and desired wavelength coverage for the NIR. Hollow cathode lamps, being significantly less expensive and easier to use, are the preferred wavelength calibration solution for most astrophysical spectrographs. We refer the reader to \\citet{kerber2007future} for a comprehensive review of the design, construction, and operation of hollow cathode lamps. Thorium ($^{232}$Th), an element often used as the cathode for such lamps, exhibits many of the desired characteristics of an atomic emission calibration source: it has many energy levels (leading to many lines), a heavy nucleus, a very long half-life, and occurs in nature as a single isotope. Uranium shares all of these characteristics except for the last: it has three naturally-occurring isotopes, the major component being uranium-238 (99.275\\%), with much smaller amounts of uranium-$235$ ($0.720\\%$) and uranium-$234$ ($0.005\\%$) \\citep{rosman1998isotopic}. These isotopes could limit the achievable precision, but the second-most abundant isotope, uranium-$235$, has hyperfine structure, so the fine structure transitions are spread over several lines and are thus much less intense.  For practical purposes, all of the uranium lines are observed as ``single'' lines in the FTS spectra. In this paper we propose an inexpensive and simple solution to the calibration problem in the NIR: the use of uranium hollow cathode lamps, which provide a high density of emission lines across the NIR spectral region.  Prior works have noted that uranium might prove a higher line density calibration source than thorium \\citep{2003JQSRT..78....1E}. Experiments with a variety of metals and fill gases in hollow cathode lamps on the Pathfinder spectrographic testbed \\citep{ramsey2008pathfinder, 2010SPIE.7735E.231R} between $1$ $\\mu$m and $1.6$ $\\mu$m confirms this. Figure~\\ref{fig:comparison} shows $25$-second exposures of Th/Ar, U/Ar, and U/Ne obtained with Pathfinder. All lamps were run at $14$ mA and share the same optical alignment. In addition to the greater line density of uranium lines, we note that neon lines in our U/Ne spectra are much less intense than argon lines in our U/Ar spectra. These results are typical for our observations in the NIR.  Using archival Fourier transform spectrometer (FTS) data from the Kitt Peak National Solar Observatory (hereafter, NSO), we have identified \\nlines uranium lines between $0.85$ $\\mu$m and $4$ $\\mu$m, most of which are below $2$ $\\mu$m. This line list complements work by \\citet{palmer1980atlas}, who provided uranium identifications for bright uranium lines from $26~000$ cm$^{-1}$ down to $11~000$ cm$^{-1}$ ($3846$ \\AA to $9091$ \\AA), and \\citet{1984ADNDT..31..299C}, who compiled a list of uranium lines between $1.8$ $\\mu$m and $5.5$ $\\mu$m. Our goal is to provide the astrophysical community with a list of largely unblended wavelength standards that can be used for the calibration of current and future NIR spectrographs.  In \\S\\ref{sec:data} we describe the archival FTS data from which we derived our measurements and line lists, and in \\S\\ref{sec:linelist} we describe the construction of the line lists, and the determination of uncertainties. In \\S\\ref{sec:results} we present the results of this line list, and in \\S\\ref{sec:discussion} we compare our line list to historical uranium line lists and to recent thorium line lists. In \\S\\ref{sec:conclusions} we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have measured \\nlines uranium lines in the NIR, which represents a significant increase in the number of calibration lines over the current NIR emission standard, Th/Ar.  This line list has been generated specifically with astronomical applications in mind --- we have taken care to avoid blended lines, and where we could not do so, we have indicated suspected or likely blends in the table.  We anticipate this line list will be useful on several upcoming astronomical spectrographs, including APOGEE \\citep{allende2008apogee}, CARMENES \\citep{2010SPIE.7735E..37Q}, and iSHELL \\citep{2008SPIE.7014E.208T}."
    },
    "1107/1107.0484_arXiv.txt": {
        "abstract": " ",
        "introduction": "The existence of very compact star clusters around the centers of many galaxies is well ascertained, nowadays. These clusters, called  Nuclear Star Clusters (NSCs) are among the densest star clusters observed, with effective radii of a few pc and central luminosities up to $10^7$ L$_\\odot$. Recently, the study of dense stellar systems has raised new interest because of the discovery, in an ever increasing number of galaxies, of compact nuclei in form of stellar resolved systems and/or NSCs which are now known to be present in galaxies across the whole Hubble sequence and not only in the dE, N galaxies \\citep{bek10}. A radial profile and velocity structure for a NSC can be determined only for the Milky Way NSC which is close enough  that it can be resolved into individual stars. The Milky Way NSC has an estimated mass of $10^7$ M$_\\odot$, and it hosts a massive black hole whose mass, $4.3\\times 10^6$ M$_\\odot$, is uniquely well determined. \\\\ The formation mechanism of nuclear star clusters  is unknown. Two competing models are possible. In the gas model, a NSC can form from the gas that migrates to the center of the galaxy where then forms stars. A variety of scenarios have been proposed to account for the required fast radial inflow of gas into the galactic center, including the magneto-rotational instability in a differentially rotating gas disk, tidal compression in shallow density profiles or dynamical instabilities. Alternatively, in the merger model, massive clusters migrate to the center via dynamical friction and merge to form a dense nucleus \\citep{tre75,capdol93}. Observations of NSCs in dE galaxies suggest that the majority, but not all, dE nuclei, could be the result of packing mass in form of orbitally decayed globular clusters (GCs) \\citep{lot04}. Numerical simulations have also shown that the basic properties of NSCs, including their shape, mass density profile, and mass-radius relation are reproduced in the merger model under a variety of explored conditions \\citep{capdol08a,capdol08b,Har11}. ",
        "conclusions": "As already shown by previous papers \\citep{capdol93,capdol08a,capdol08b} the dynamical friction dragging of massive globular clusters toward the galactic center is a viable explanation of the local high densities. The novelty of this work is the inclusion in the simulations of a massive black hole in the galactic center. Its role, although important in the vicinity of its influence radius, does not alter the general characteristics of the merger event of a set of GCs falling to the center. The merger reaches a sort of quasi-steady state, slowly evolving due to internal relaxation which is partly affected by the presence of the massive black hole. With regard to the Milky Way nuclear cluster, it is relevant noting how the original `core' profile of the GCs (the NSC building blocks) is maintained for a time long enough to justify the core actually observed in the MW NSC. Moreover, the outward profile radial slope, $\\rho \\propto r^{-2}$, of the merger system is in good agreement (within 30 pc from the center) with the observed  MW NSC profile. Surely, an even more stringent answer to the question of the MW NC formation will come from a high resolution, well extended in time numerical study of the secular evolution of the merger remnant as well as from a detailed study of the expected observed stellar population (both in progress). \\small  %"
    },
    "1107/1107.2397_arXiv.txt": {
        "abstract": "Deep $K$-band imaging of the most luminous $z \\simeq 4$ quasars currently offers the earliest possible view of the mass-dominant stellar populations of the host galaxies which house the first super-massive black holes in the Universe. This is because, until the advent of the James Webb Space Telescope, it is not possible to obtain the necessary deep, sub-arcsec resolution imaging at rest-frame wavelengths $\\lambda_{rest} > 4000$\\AA\\ at any higher redshift. We here present and analyse the deepest, high-quality $K_S$-band images ever obtained of luminous quasars at $z \\simeq 4$, in an attempt to determine the basic properties of their host galaxies less than 1 Gyr after the first recorded appearance of black holes with $M_{bh} > 10^9\\,{\\rm M_{\\odot}}$. To maximise the robustness of our results we have carefully selected two Sloan Digital Sky Survey quasars at $z \\simeq 4$. With absolute magnitudes $M_i < -28$, these quasars are representative of the most luminous quasars known at this epoch but they also, crucially, lie within 40 arcsec of comparably-bright foreground stars (required for accurate point-spread-function definition), and have redshifts which ensure line-free $K_S$-band imaging. The data were obtained in excellent seeing conditions ($<0.4$-arcsec) at the European Southern Observatory on the Very Large Telescope with integration times of $\\simeq 5.5$ hours per source. Via carefully-controlled separation of host-galaxy and nuclear light, we estimate the luminosities and stellar masses of the host galaxies, and set constraints on their half-light radii. The apparent $K_S$-band magnitudes of the quasar host galaxies are consistent with those of luminous radio galaxies at comparable redshifts, suggesting that these quasar hosts are also among the most massive galaxies in existence at this epoch. However, the quasar hosts are a factor $\\sim 5$ smaller ($\\langle r_{1/2} \\rangle = 1.8\\,{\\rm kpc}$) than the host galaxies of luminous low-redshift quasars. We estimate the stellar masses of the $z \\simeq 4$ host galaxies to lie in the range $2 - 10 \\times 10^{11}\\,{\\rm M_{\\odot}}$, and use the CIV emission line in the Sloan optical spectra to estimate the masses of their central supermassive black holes. The results imply a black-hole:host-galaxy mass ratio $M_{bh}:M_{gal} \\simeq 0.01 - 0.05$. This is an order of magnitude higher than typically seen in the low-redshift Universe, and is consistent with existing evidence for a systematic growth in this mass ratio with increasing redshift (i.e. $M_{bh}:M_{gal} \\propto (1+z)^{1.4-2.0}$), at least for objects selected as powerful active galactic nuclei. ",
        "introduction": "At all but the most modest redshifts, the measurement of the velocity width of broad permitted lines (e.g. H$\\beta$, MgII, CIV) from unobscured Type-1 active galactic nuclei (AGN) offers the only direct way to estimate the masses of super-massive black holes (Wandel, Peterson \\& Malkan 1999; Kaspi et al. 2000; McLure \\& Dunlop 2001; Ho 2002; McLure \\& Jarvis 2002; Willott et al. 2003; McLure \\& Dunlop 2004; Vestergaard \\& Peterson 2006; Peng et al. 2006a; Vestergaard et al. 2008; Willott et al. 2010). For this reason, determining the properties of AGN host galaxies remains of crucial importance for exploring the relationship between black-hole and galaxy evolution over cosmic history, and in particular for testing alternative theories for the origin of the now well-established proportionality between black-hole and galaxy bulge mass found in the local/present-day Universe (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Silk \\& Rees 1998; Gebhardt et al. 2000; Merritt \\& Ferrarese 2001; McLure \\& Dunlop 2002; Tremaine et al. 2002; Bettoni et al. 2003; Marconi \\& Hunt 2003, Merloni et al. 2010; Jahnke \\& Maccio 2011).  Quasars, as the most extreme of AGN, offer potentially the greatest challenge to theory, especially at high redshift where, for example, black holes with masses $M_{bh} > 1 \\times 10^{9}\\,{\\rm M_{\\odot}}$ already appear to have been in place less than 1 billion years after the Big Bang (Fan et al. 2001, 2003; Willott et al. 2007; Mortlock et al. 2011). However, extracting reliable information on the host galaxies of these overwhelmingly bright AGN presents a formidable technical challenge.  At low redshift ($z < 0.4$) this challenge was effectively met when the first refurbishment of the {\\it Hubble Space Telescope} ({\\it HST}) enabled optical imaging of the necessary depth {\\it and} exquisite+stable angular resolution to allow the scalelengths and morphologies of low-redshift quasar hosts to be reliably determined for the first time (e.g. Disney et al. 1995; McLure et al. 1999; Dunlop et al. 2003; Floyd et al. 2004). Indeed, the results of these deep {\\it HST} optical imaging programs of quasar hosts, when combined with ``virial'' measurements of the black-hole masses in the same objects, were instrumental in demonstrating that these most active objects displayed the same black-hole:host-galaxy mass proportionality as found in quiescent galaxies (McLure \\& Dunlop 2002).  However, extending the effective study of quasar host galaxies out to higher redshifts, and in particular into the ``quasar epoch'' at $z > 2$, has proved to be extremely difficult. There are a number of reasons for this. First, there is the obvious cosmological dimming of the host-galaxy light. Second there is the necessity of observing in the near-infrared, in order to continue to sample the rest-frame optical light of the host galaxy (both for comparison with low-redshift studies, and to avoid the host-galaxy light being completely swamped by the UV-bright quasar). Some success in detecting the hosts of quasars out to $z \\simeq 2$ was achieved with the NICMOS infrared camera on {\\it HST} (Kukula et al. 2001, Ridgway et al. 2001; Peng et al. 2006b) but, even with this high-resolution near-infrared imaging, obtaining reliable measurements of host-galaxy luminosities and scalelengths proved extremely difficult. We now know that this was almost certainly in part due to the fact that, in general, massive galaxies at high-redshift have since been discovered to be generally much more compact than their low-redshift counterparts (Daddi et al. 2005; Trujillo et al. 2006, 2007; Longhetti et al. 2007; Zirm et al. 2007; Cimatti et al. 2008; van Dokkum et al. 2008; Buitrago et al. 2008; Szomoru et al. 2010; Targett et al. 2011). If this is also true for quasar hosts, then the separation of galaxy and nuclear light in quasars at $z \\ge 2$ will inevitably be even more problematic than was perhaps first anticipated.  Nonetheless, these difficulties have not deterred several groups from attempting to repeat this experiment at redshifts as high as $z \\simeq 4 - 5$. Such studies have been encouraged (at least in part) by the advent of active/adaptive optics on ground-based 8-m class telescopes, and by the fact that $z \\simeq 4$ is the natural limit for ground-based studies of quasar host galaxies if one hopes to detect host-galaxy light longward of the 4000\\AA/Balmer break in the $K$-band.  However, perhaps unsurprisingly, existing studies of the host galaxies of quasars at $z \\simeq 4$ have met with very limited success. For example, using {\\it HST} NICMOS $H$-band data, Peng et al. (2006b) attempted to measure the luminosities and sizes of the hosts of two gravitationally-lensed quasars at $z=4.1$ and $z=4.5$ via two-dimensional modelling. Although reliable host-galaxy luminosities were claimed, the {\\it HST} images sampled light shortward of $\\lambda_{rest} = 4000$\\,\\AA\\ and, in addition, the complexity of the lensing model led the authors to conclude that the extracted morphological parameters could not be trusted. In an alternative ground-based approach, Hutchings (2003, 2005) observed seven quasars at $z \\simeq 5$ with the Gemini 8-m telescope in the $J$, $H$, and $K$-bands. Using point-spread function (PSF) subtraction, host-galaxy luminosities were estimated from the PSF-subtracted images, although no evidence of smooth centrally-concentrated host galaxies was found and, again, this time due to the decision to observe quasars at $z \\simeq 5$, this imaging sampled rest-frame wavelengths $\\lambda_{rest} < 4000$\\,\\AA, limiting the usefulness of the extracted luminosities. Most recently, McLeod \\& Bechtold (2009) observed 34 $z \\simeq 4$ quasars in the $K$-band with the Magellan I and Gemini North telescopes. However, despite the use of well-controlled PSF subtraction, only four host galaxies were detected in their deepest and sharpest data, at an average quasar redshift $\\langle z \\rangle = 3.8$. Host-galaxy luminosity and size estimates were attempted from the PSF-subtracted images of these 4 quasars, although the authors noted that the surface-brightness limit of these data was really of inadequate depth for the accurate recovery of host-galaxy sizes. The deepest imaging obtained by these authors ($\\simeq 3.5$ hours per source on 8-m class telescopes) would thus seem to represent a minimum requirement for the discovery and study of quasar host galaxies at $z \\simeq 4$.  In this paper we present and analyse even deeper ($\\simeq 5.5$ hours per source) European Southern Observatory (ESO) 8.2-m Very Large Telescope (VLT) $K_S$-band imaging of two $z \\simeq 4$ Sloan Digital Sky Survey (SDSS) quasars. These two quasars were carefully chosen to have a redshift which ensures line-free imaging of the host-galaxy stellar population at $\\lambda_{rest} > 4000$\\,\\AA, and to have a star of comparable brightness within an angular radius of 40 arcsec, in order to provide a robust, high signal:noise ratio, on-detector representation of the VLT $K_S$-band PSF over the complete duration of the host-galaxy imaging. This care in quasar target selection, combined with deliberate (flexibly scheduled) use of only the very best seeing conditions (and the deliberate avoidance of adaptive optics), has enabled us to achieve robust detections of the host galaxies of both quasars, despite the fact that, with $M_i < -28$, these objects are representative of the most luminous/massive black holes in existence at this (or indeed any) epoch.  This paper is structured as follows. In Section 2 we describe our observing strategy, our target selection, and summarise the new near-infrared data obtained with the VLT. Next, in Section 3 we explain the PSF subtraction technique used to separate host and nuclear light. Then, in Section 4 we describe the determination of host-galaxy properties and the estimation of the central black-hole masses. Finally, in Section 5 we attempt to place our results in the context of studies of other galaxy populations at both high and low redshift, and briefly explore the consequences of this study for the inferred cosmic evolution of the black hole-to-galaxy mass ratio. Throughout we adopt a cosmology with $H_{0}=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{m}=0.3$, and $\\Omega_{\\Lambda}=0.7$. Unless otherwise noted, we report all magnitudes in the Vega system to ease comparison with previous studies, but our $K_S$-band photometry can be simply converted to the AB magnitude system via $K_{S,AB} = K_{S,Vega} + 1.85$. ",
        "conclusions": "We have obtained and analysed the deepest, high-quality $K_S$-band images ever obtained of luminous quasars at $z \\simeq 4$, in an attempt to determine the basic properties of their host galaxies less than 1 Gyr after the first recorded appearance of black holes with $M_{bh}~>~10^9\\,{\\rm M_{\\odot}}$. To maximise the robustness of our results, we carefully selected two SDSS quasars at $z \\simeq 4$. With absolute magnitudes $M_i < -28$, these quasars are representative of the most luminous quasars known at this epoch but they also, crucially, lie within 40 arcsec of comparably-bright foreground stars (required for accurate point-spread-function definition), and have redshifts which ensure line-free $K_S$-band imaging.  Our new data were obtained in excellent seeing conditions ($<0.4$-arcsec) with ISAAC on the ESO VLT, with integration times of $\\simeq 5.5$ hours per source. Via carefully-controlled separation of host-galaxy and nuclear light, followed by two-dimensional galaxy model-fitting, we have estimated the luminosities and stellar masses of the host galaxies. In addition, although it did not prove possible to determine basic morphological type (e.g. S\\'{e}rsic index $n$), we have extracted what appear to be robust estimates of the host-galaxy half-light radii.  We find that the $K_S$-band magnitudes of the quasar host galaxies are consistent with those of luminous radio galaxies at comparable redshifts. This suggests that these quasar hosts are also among the most massive galaxies in existence at this epoch, and indeed inferred stellar masses of these host galaxies lie in the range $2 - 10 \\times 10^{11}\\,{\\rm M_{\\odot}}$, depending on assumed formation redshift. We also find that these massive $z \\simeq 4$ quasar host galaxies are a factor $\\sim 5$ smaller ($\\langle r_{1/2} \\rangle = 1.8\\,{\\rm kpc}$) than the host galaxies of luminous low-redshift quasars, or indeed quiescent galaxies of comparable stellar mass. However, we argue that this is as expected given the growing evidence for the increasing compactness of massive galaxies with increasing redshift, and conclude that, just as at low redshift, the properties of quasar host galaxies at $z \\simeq 4$ are broadly as expected for ``normal'' galaxies of comparable mass at the relevant cosmological epoch.  Finally we have used the CIV emission line in the SDSS optical spectra of the quasars to estimate the masses of their central supermassive black holes. The results imply extreme black-hole masses approaching $M_{bh} \\simeq 10^{10}\\,{\\rm M_{\\odot}}$, but we argue that this is not unrealistic given the extreme nature and rarity of our target $z \\simeq 4$ SDSS quasars (i.e. $1 - 2$ per comoving Gpc$^3$), and current observational and theoretical understanding of black-hole growth in the young Universe. Combining the black-hole and host-galaxy mass measurements we infer a black-hole:host-galaxy mass ratio $M_{bh}:M_{gal} \\simeq 0.01 - 0.05$. This is an order of magnitude higher than typically seen in the low-redshift Universe, and is consistent with existing evidence for a systematic growth in this mass ratio with increasing redshift, at least for objects selected as powerful active galactic nuclei."
    },
    "1107/1107.1661.txt": {
        "abstract": "In this paper we review the story of the \\sax\\ discovery of the Gamma Ray Burst afterglow and their cosmological distance, starting from their first detection with Vela satellites and from the efforts done before \\sax. We also discuss the consequences of the \\sax\\ discovery, the issues left open by \\sax, the progress done up to now and its perspectives. ",
        "introduction": "Gamma Ray Bursts (GRBs) are the most fascinating events in the Universe. They are sudden bright flashes of gamma--ray radiation of celestial origin, with variable duration from milliseconds to several hundreds of seconds. In rare cases time duration up to thousands of seconds has been observed. Most of the emission is made of photons with energies from several keV to tens of MeV and above. With the current instrumentation their observed occurrence rate is 2--3 per day over the entire sky. Their arrival time is impredictable as it is impredictable their arrival direction. When they are on, their brightness overcomes any other celestial gamma-ray source. They were discovered by chance at the end the 60\u0092s with the American Vela spy satellites, devoted  to monitor the compliance of the Soviet Union and other nuclear-capable states with the 1963 \"Partial Test Ban Treaty\". The American militaries kept this discovery in their drawers until they decided to talk about that with scientists colleagues in Los Alamos laboratories, who published their discovery in 1973 \\cite{Klebesadel73} with the title \"Observations of Gamma-Ray Bursts of Cosmic Origin\". Once discovered, the first questions were: which are their progenitors? Are they normal stars or compact stars, like white dwarfs (WD), neutron stars (NS) or Black Holes (BH)? or more simply they are something that occurs in the interstellar medium? or in comets? At which distance are their sites? Are local sites? Are their sites in our Galaxy or in extragalactic objects? Which is the power involved?  The solution of these issues implied  complex observational problems, like the accurate localization of the events. This was  a tough task in the gamma--ray energy band, where detectors could give only coarse source directions. A possible solution of this problem could be the search of GRB counterparts at longer wavelengths, that could be a new born source or an already existing one.  Many satellite missions (e.g., the Russian satellites Venera 11, Venera 12, Prognoz 6, Prognoz 9, Konus, Granat, and the American Pioneer--Venus Orbiter and Solar Maximum Mission) which included instrumentation devoted to the detection of GRBs (see, e.g., Refs.~\\cite{Niel76,Barat81}) were performed, but a very small progress in the GRB origin was obtained (see, e.g., Ref.~\\cite{Epstein88} and the very complete review by Vedrenne and Atteia \\cite{Vedrenne09}). Both their spatial distribution and their distance was a persistent  mystery. Their origin was reminiscent of the great debate in 1920s about the local or extragalactic origin of spiral nebulae. The localizations were very coarse and no counterpart at longer wavelengths was found (see, e.g., Ref.~\\cite{Hudec87}). The most accurate localizations, with error boxes even of a few square arcminutes, were obtained using the triangulation method. This consists in the accurate timing of GRBs with omnidirectional detectors aboard at least three satellites. From the time delay in the arrival time of the same event by these satellites it is possible to derive the arrival direction of the event. However these localizations were obtained after long times, from months to years, from the GRB occurrence (e.g., Ref.~\\cite{Atteia87}).  Many theoretical models on  GRB progenitors were worked out, with the largest consensus being obtained by the models which assumed that GRBs originate in Galactic disk neutron stars.  The largest observational effort  was performed with the BATSE experiment  aboard the NASA \\gro\\ satellite (CGRO) \\cite{Fishman94}, with a very wide field of view (2$\\pi$) (see below). In spite of the great wealth of information, the sites of GRBs were not discovered. The real revolution came with the \\sax\\ satellite that, in only 6 months from its launch (30 April 1996), allowed to discover the X--ray counterparts of GRBs and their distance. After about 30 years from their discovery, we eventually learnt that GRBs are huge explosions in galaxies at cosmological distances.  In this paper we review the major results obtained in the pre--\\sax\\ era, the \\sax\\ revolution and its major results obtained. We will also discuss the consequences of the \\sax\\ discovery, the issues left open  and the further  progress obtained with later missions devoted to GRBs. Finally we will discuss some relevant still open issues and the prospects to solve them. ",
        "conclusions": ""
    },
    "1107/1107.4787_arXiv.txt": {
        "abstract": "We present some theoretical predictions concerning the amplitude and magnetic sensitivity of the linear polarization signals produced by scattering processes in the hydrogen Ly${\\alpha}$ line of the solar transition region. To this end, we have calculated the atomic level polarization (population imbalances and quantum coherences) induced by anisotropic radiation pumping in semi-empirical and hydrodynamical models of the solar atmosphere, taking into account radiative transfer and the Hanle effect caused by the presence of organized and random magnetic fields. The line-center amplitudes of the emergent linear polarization signals are found to vary typically between 0.1\\% and 1\\%, depending on the scattering geometry and the strength and orientation of the magnetic field. The results shown here encourage the development of UV polarimeters for sounding rockets and space telescopes with the aim of opening up a diagnostic window for magnetic field measurements in the upper chromosphere and transition region of the Sun. ",
        "introduction": "One of the big challenges of 21st century astrophysics is to understand the magnetism of the Sun, and in so doing to develop the tools needed for the exploration of the magnetic activity in other types of stars across the Hertzsprung-Russel diagram. A crucial step toward this goal is to decipher the magnetic structure of the upper chromosphere, the key interface region between the underlying (cooler) photosphere and the overlying (hotter) corona where dominance of the physics passes from hydrodynamic to magnetic forces (e.g., the review by Harvey 2006). To this end, we need:  \\begin{itemize}  \\item to identify observables sensitive to the magnetic fields of the upper chromosphere and transition region;  \\item to develop suitable diagnostic tools to infer the magnetic field from those observables;  \\item to design and build the telescopes and instrumentation needed for measuring the observables.  \\end{itemize}   The upper chromosphere is indeed a very important part of the Sun because all of the plasma, magnetic field and energy in the corona and solar wind are supplied through this key boundary region. The observables that contain information on the physical conditions of this region are mainly spectral lines in the ultraviolet (UV) and far-UV (FUV) spectral range, with Ly${\\alpha}$ being the most prominent line. The intensity profile (i.e., the Stokes $I(\\lambda)$ profile) of Ly${\\alpha}$ and other lines of the Lyman series have been measured on the solar disk by several instruments on board rockets and space-based telescopes (e.g., Roussel-Dupr\\'e 1982; Warren et al. 1998). These measurements show that when observed on the solar disk the hydrogen Lyman lines are always in emission, and that the emission originates in the upper chromosphere and transition region. Interestingly, high-resolution Ly${\\alpha}$ {\\em intensity} images of quiet-Sun regions taken with the Very high Angular resolution ULtraviolet Telescope (VAULT) sounding rocket show a forest of elongated fibrils suggesting closed magnetic loops in the upper chromosphere (Vourlidas et al. 2010). However, no quantitative information on the magnetic fied of the upper solar atmosphere can be obtained from this type of observations because the intensity profile of Ly${\\alpha}$ and of many other transition region lines is practically insensitive to the strength, inclination and azimuth of the magnetic field vector.  The only way to obtain quantitative empirical information on the strength and orientation of the magnetic field is through the measurement and interpretation of the polarization that some physical mechanisms introduce in spectral lines. Unfortunately, the familiar Zeeman effect is of little practical interest for the ``measurement\" of magnetic fields in the upper chromosphere and transition region (except perhaps in sunspots) because the circular polarization (Stokes $V$) scales with the ratio, ${\\cal R}$, between the Zeeman splitting and the Doppler line-width and the linear polarization (Stokes $Q$ and $U$) scales with ${\\cal R}^2$; this ratio is very small for the UV and FUV lines that originate in the weakly magnetized plasma of the upper solar chromosphere because ${\\cal R}\\,{\\propto}\\,{\\lambda}{B}/{\\sqrt{T}}$ (with $\\lambda$ the wavelength, $B$ the magnetic strength and $T$ the kinetic temperature; see Landi Degl'Innocenti \\& Landolfi 2004). Fortunately, there is yet another physical mechanism by means of which the magnetic fields of a stellar atmosphere leave fingerprints in the spectral line polarization: the Hanle effect. The absorption of {\\em anisotropic} radiation produces atomic level alignment (i.e., the individual magnetic substates of levels with angular momentum $j{\\ge}1$ are unevenly populated, in such a way that the populations of sublevels with different values of $|M|$ are different and, moreover, quantum coherences between them may appear). This, in turn, gives rise to linear polarization in the spectral line under consideration (e.g., Trujillo Bueno 2001). The Hanle effect due to the presence of a magnetic field inclined with respect to the symmetry axis of the pumping radiation field modifies the atomic level alignment and the linear polarization of the spectral line radiation. Approximately, the emergent linear polarization is sensitive to magnetic strengths between 0.2$B_H$ and 5$B_H$, where $B_H{=}\\,{1.137{\\times}10^{-7}}/{t_{\\rm life}g}$ is the critical Hanle field (in G) for which the Zeeman splitting of the line's level under consideration is similar to its natural width (with $t_{\\rm life}$ the level's radiative lifetime, in seconds, and $g$ its Land\\'e factor). Note that $B_H{\\approx}50$ G for Ly${\\alpha}$, $B_H{\\approx}20$ G for Ly${\\beta}$ and $B_H{\\approx}8$ G for Ly${\\gamma}$, and that the magnetic sensitivity provided by the Hanle effect is independent of the Doppler line width.  Each Lyman line results from two blended transitions between a lower level, with $j=1/2$, and two upper levels, with $j=1/2$ and $j=3/2$. Therefore, for each Lyman line the only level that can be aligned and contribute to the emergent linear polarization is the upper level with $j=3/2$ (when the hyperfine structure of hydrogen is neglected, which is known to be a good approximation for Ly${\\alpha}$). A particularly important question is whether the atomic alignment of the $2p^2{\\rm P}_{3/2}$ level of Ly${\\alpha}$ is significant or, in other words, if the Ly${\\alpha}$ radiation that illuminates the hydrogen atoms of the solar atmosphere is sufficiently anisotropic in the line formation region. Obviously, the anisotropy of the incident Ly${\\alpha}$ radiation is very significant in the optically thin case of the solar corona observed off-the-limb at large distances above it (e.g., Trujillo Bueno et al. 2005; and more references therein). The key question for us here is however whether the Ly${\\alpha}$ radiation is sufficiently anisotropic within the solar transition region itself, so as to produce there a significant amount of atomic alignment in the $2p^2{\\rm P}_{3/2}$ level. As mentioned above, the Lyman radiation that we observe when pointing to the solar disk is produced by the plasma of the solar transition region, which is very optically thick at the central wavelength of Ly${\\alpha}$. As a matter of fact, the center-to-limb variation (CLV) of the observed Ly${\\alpha}$ intensity is very small or negligible (e.g., Roussel-Dupr\\'e 1982; Curdt et al. 2008), which implies that the intensity of the {\\em outgoing} radiation (i.e., that propagating outwards) shows no significant variation with the heliocentric angle $\\theta$. However, this does not mean that the illumination of the transition region atoms is isotropic -- which would imply that the scattering polarization in Ly${\\alpha}$ is zero. As we shall see below, the intensity of the {\\em incoming} radiation (i.e., that propagating inwards) calculated in semi-empirical and hydrodynamical models of the solar chromosphere shows significant CLV, so that the hydrogen atoms of the solar transition region must be illuminated by an anisotropic radiation field. The main goal of this Letter is to show that the linear-polarization amplitude of the emergent Ly${\\alpha}$ radiation is predicted to be significant because of the anisotropy of the incoming radiation and that via the Hanle effect it is sensitive to the magnetic fields expected for the upper chromosphere and transition region of the Sun.  We focus here on Ly${\\alpha}$ because it is the most intense line of the solar transition region and also because its polarization is practically insensitive to collisional depolarization there, but we point out that linear polarization due to atomic level alignment is also expected for at least Ly${\\beta}$ and Ly${\\gamma}$. ",
        "conclusions": "Our radiative transfer investigation in semi-empirical and hydrodynamical models of the solar atmosphere indicate that the Ly${\\alpha}$ line should show measurable linear polarization signals when observing the solar disk. Via the Hanle effect the line-center amplitudes turn out to be sensitive to the presence of magnetic fields in the solar transition region, with good sensitivity to field strengths between 10 and 100 G. For such field strengths there is no significant impact of the Zeeman effect on the Ly${\\alpha}$ radiation.  The above-mentioned prediction is based on the same quantum theory that explains the linear polarization profiles observed in various visible and near-IR lines (e.g., Manso Sainz \\& Trujillo Bueno 2003; \\v{S}t\\v{e}p\\'an \\& Trujillo Bueno 2010). This CRD theory is suitable also for estimating the line-center polarization amplitudes of strong resonance lines like Ly${\\alpha}$, but we point out and advance here that partial frequency redistribution (PRD) effects do produce interesting and sizable $Q/I$ signals in the Ly${\\alpha}$ wings (qualitatively similar to those shown in figure 9 of Trujillo Bueno 2011 for the Mg {\\sc ii} k-line at 2795 \\AA).  The predicted line-center amplitudes of the Ly${\\alpha}$ scattering polarization are smaller than 1\\%. Nevertheless, with a moderate aperture telescope they can be measured with a polarimetric sensitivity of 0.1\\% and a spectral resolution of 0.1 \\AA\\ by opting for a spatial resolution of the order of 1 arcsecond and a temporal resolution of the order of 1 minute (e.g., see the Ly${\\alpha}$ polarization sounding rocket project outlined in Ishikawa et al. 2011). Information on the thermal and magnetic structure of the solar transition region can be obtained from the observed $Q/I$ and $U/I$ profiles themselves and through detailed radiative transfer simulations in given atmospheric models. Therefore, it would be worthwhile to develop a spectropolarimeter for measuring from space the linear polarization profiles caused by scattering processes and the Hanle effect in the Ly${\\alpha}$ line of the solar transition region. Unfortunately, the pioneering Soviet satellite experiment of Stenflo et al. (1980) failed because the optical transmission of the instrument was severely degraded due to in-orbit contamination from the satellite. Such a development would open up a novel diagnostic window for the exploration of the magnetism of the outer solar atmosphere.    {\\bf Acknowledgments} We are grateful to S. Tsuneta (NAOJ), K. Kobayashi (UAH) and all the other members of the Chromospheric Ly${\\alpha}$ Spectropolarimeter (CLASP) sounding rocket project for their interest in the theoretical results presented here and for stimulating scientific discussions. Thanks are also due to R. Manso Sainz (IAC) and L. Belluzzi (IAC) for their careful review of this letter, and to M. Carlsson (UiO) for kindly providing the hydrodynamical model atmosphere. Financial support by the Spanish Ministry of Science through projects AYA2010-18029 (Solar Magnetism and Astrophysical Spectropolarimetry) and CONSOLIDER INGENIO CSD2009-00038 (Molecular Astrophysics: The Herschel and Alma Era) is gratefully acknowledged."
    },
    "1107/1107.4264_arXiv.txt": {
        "abstract": "We present a highly parallel implementation of the cross-correlation of time-series data using graphics processing units (GPUs), which is scalable to hundreds of independent inputs and suitable for the processing of signals from ``Large-\\(N\\)'' arrays of many radio antennas.  The computational part of the algorithm, the X-engine, is implementated efficiently on Nvidia's Fermi architecture, sustaining up to 79\\% of the peak single precision floating-point throughput. We compare performance obtained for hardware- and software-managed caches, observing significantly better performance for the latter. The high performance reported involves use of a multi-level data tiling strategy in memory and use of a pipelined algorithm with simultaneous computation and transfer of data from host to device memory.  The speed of code development, flexibility, and low cost of the GPU implementations compared to ASIC and FPGA implementations have the potential to greatly shorten the cycle of correlator development and deployment, for cases where some power consumption penalty can be tolerated. ",
        "introduction": "\\label{sec:introduction}  We apply graphics processing units (GPUs) to the problem of signal processing for radio astronomy.  While not a classic high performance computing (HPC) application, there are now many radio astronomy applications that require in excess of O(100) tera-floating-point operations per second (TFLOPS) sustained performance.  The Square Kilometer Array (SKA) due to be built circa 2020, will raise the processing needs into the exascale regime~\\cite{SKA, cornwell-exaflop}.   A typical processing pipeline consists of: digitization of the raw voltage time-series from individual antennas; cross-correlatation; instrument electronics calibration and Fourier imaging reconstruction of the sky.  The extreme computational cost lies predominantly in the cross-correlation stage; this requires that the signal from every antenna is correlated with every other, and scales quadratically with the number of antenna.\\footnote{Alternative approaches that scale as \\(N \\log N\\) have been proposed~\\cite{PhysRevD.82.103501}, but cross-correlation, thus far the mainstay technique, enables greater generality in data calibration and range of scientific application.} Science applications that demand high dynamic range and sensitivity in imaging drive interest in arrays of hundreds to tens of thousands of collectors (e.g., antennas).  This raises the processing needs from being modest and easily manageable, into the HPC domain, e.g., the Murchison Widefield Array, currently in prototype stage in Australia will require O(10-100) TFLOPS sustained for cross-correlation of configurations from 128 to 512 collectors~\\cite{Lonsdale:2009cb}.    The application of GPUs to cross-correlation is not uncharted territory: there have been several works investigating GPUs use for this very purpose~\\cite{Harris:2008, Wayth:2009, Nieuwpoort09usingmany-core}.  While these demonstrated that GPUs are suitable for this task, in all cases only 10--30\\% of the GPU's peak performance was obtained and the problem was described as being {\\it bandwidth bound}.  This work presents an approach to cross-correlation that is catered to the deep memory hierarchy of Nvidia's Fermi GPU.  We compare the performance obtainable using both hardware- and software-managed caches, the latter of which is more familiar to GPU programmers.  We find in favor of the software-managed cache, achieving up to 79\\% of peak performance, equating to performance in excess of 1 TFLOPS on the GeForce GTX 480.  More significantly, our approach is scalable to future architectures which will likely feature a greater disparity between compute throughput and memory bandwidth.   This paper is split up as follows; in \\S\\ref{sec:cross-correlation} we describe the cross-correlation process, specifically the FX algorithm. An overview of GPUs and previous attempts at using GPUs for this problem is given in \\S\\ref{sec:gpus}.  We describe our kernels in \\S\\ref{sec:implementation}, where we contrast the hardware- and software-managed cache implementations.  In \\S\\ref{sec:pcie} we consider performance of the integrated system where we include the overhead of PCIe transfers.  We discuss the implications of our results in \\S\\ref{sec:discuss} before concluding in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have presented an implementation of the cross-correlation algorithm using Nvidia's Fermi architecture.  This application is very well matched to the architecture, and sustains in excess of 1 TFLOPS for a large number of stations (79\\% of theoretical peak).  Key to obtaining this performance was the use of a software-managed cache and a multi-level tiling strategy.  This performance can be sustained when streaming the input data over the PCIe bus.  This represents a significant improvement over previous work on GPUs and comparison to other commodity platforms.  While not power competitive compared to FPGA solutions, its increased flexibility and lower development costs make it an attractive solution.  With a fully optimized GPU X-engine implemented, the next step is to develop this into a full FX correlator.  The design we are pursuing is a hybrid approach, utilizing an FPGA solution for the F-engine and likely a packetized switch approach for the corner turn~\\cite{2008PASP..120.1207P}.  Beyond this, future work will include increasing the efficiency of the X-engine in the small \\(N\\) regime.  Although the work here primarily targetted the large \\(N\\) regime for which there are significant science drivers~\\cite{Lonsdale:2009cb}, most current radio telescope array installations are of the small \\(N\\) type.  By maximizing the efficiency at low \\(N\\), e.g., 32 stations, this would increase the applicability of GPU computing for signal correlation today, giving a viable alternative to the usual FPGA solutions."
    },
    "1107/1107.3358_arXiv.txt": {
        "abstract": "The HII region IC 1848 harbors a lot of intricate elephant trunk-like structures that look morphologically different from usual bright-rimmed clouds (BRCs). Of particular interest is a concentration of thin and long elephant trunk-like structures in the southeastern part of IC 1848E. Some of them have an apparently associated star (or two stars) at their very tip. We conducted $VI_{c}$ photometry of several of these stars. Their positions on the $V/(V-I_{c})$ color-magnitude diagram as well as the physical parameters obtained by SED fittings indicate that they are low-mass pre-main-sequence stars having ages of mostly one {\\it Myr} or less. This strongly suggests that they formed from elongated, elephant trunk-like clouds. We presume that such elephant trunk-like structures are genetically different from BRCs, on the basis of the differences in morphology, size distributions, and the ages of the associated young stars. We suspect that those clouds have been caused by hydrodynamical instability of the ionization/shock front of the expanding HII region. Similar structures often show up in recent numerical simulations of the evolution of HII regions. We further hypothesize that this mechanism makes a third mode of triggered star formation associated with HII regions, in addition to the two known mechanisms, i.e., {\\it collect-and-collapse} of the shell accumulated around an expanding HII region and {\\it radiation-driven implosion} of BRCs originated from pre-existing cloud clumps. ",
        "introduction": "Recent high-resolution images of many HII regions taken with the {\\it Hubble Space Telescope\\/} and the {\\it Spitzer Space Telescope} show very complicated structures inside them. One of such HII regions is IC 1848 (= W5). See, e.g., Figure 4 of Koenig et al. (2008), where we find a wealth of intricate structures inside/on its boundaries. Some of them are bright-rimmed clouds (BRCs) cataloged in Sugitani, Fukui, and Ogura (1991). But others are morphologically much different from usual BRCs, suggesting that they are genetically different from BRCs. We discuss this point in Sect. 4, but briefly, we suspect that, whereas BRCs mostly originate from pre-existing cloud clumps left-over in evolved HII regions, some of them may have resulted from the hydrodynamical instability of the ionization front of the expanding HII region. Of particular interest is a concentration of thin and long elephant trunk-like structures (hereafter abbreviated as {\\it ETLS\\/}s) in the southeastern part of IC 1848E. Figure 1 is a contrast-enhanced pseudo-color image of part of IC 1848E taken by the {\\it Spitzer Space Telescope\\/} (blue : 3.6$\\mu$m, green : 8.0$\\mu$m, red : 24.0$\\mu$m). Note that all these {\\it ETLS\\/}s point to HD 18326, the exciting O star of IC 1848E. Very interestingly, some of them have a star/a few stars at their very tip, as marked in figure 1. This led us to suspect that they gave birth to these stars under the compressing effects of HII gas. Zavagno et al. (2007) found two intrusions of similar morphology with a star at their tip in RCW 120 (see their Fig. 12). We further suspect that the hydrodynamical instability of the ionization fronts creating {\\it ETLS\\/}s makes a third mechanism of triggered star formation associated with HII regions, in addition to the {\\it collect-and-collapse} process of the shell accumulated around an expanding HII region and {\\it radiation-driven implosion} of BRCs.  In order to examine the pre-main-sequence (PMS) nature of the stars located at the tip of the {\\it ETLS\\/}s, we carried out $VI_{c}$ photometry of these stars and constructed a $V/V-I_{c}$ color-magnitude diagram (CMD). We also used near-infrared (NIR) data from the {\\it Two Micron All Sky Survey\\/} (2MASS) to construct a NIR color-color diagram as well as mid-infrared (MIR) data from the {\\it Spitzer Space Telescope} to make spectral energy distribution (SED) curves.  \\begin{figure} \\begin{center} \\includegraphics[scale = 1.5, trim = 0 0 0 0, clip]{Fig1.eps} \\end{center} \\caption{Contrast-enhanced {\\it Spitzer} pseudo-color image of part of IC 1848E taken from the NASA {\\it Spitzer Space Telescope} website. Stars at the tip of elephant trunk-like structures are marked together with two bright-rimmed clouds and the exciting star of IC 1848E. The scale is shown. North is up, east to the left.} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "1107/1107.1488_arXiv.txt": {
        "abstract": "This study is the first in a series in which we analyze the structure and topology of the Cosmic Web as traced by the Sloan Digital Sky Survey.  The main issue addressed in the present study is the translation of the irregularly distributed discrete spatial data in the galaxy redshift survey into a representative density field. The density field will form the basis for a statistical, topological and cosmographic study of the cosmic density field in our Local Universe.  We investigate the ability of three reconstruction techniques to analyze and investigate weblike features and geometries in a discrete distribution of objects. The three methods are the linear Delaunay Tessellation Field Estimator (DTFE), its higher order equivalent Natural Neighbour Field Estimator (NNFE) and a version of Kriging interpolation adapted to the specific circumstances encountered in galaxy redshift surveys, the {\\it Natural Lognormal Kriging} technique. DTFE and NNFE are based on the local geometry defined by the Voronoi and Delaunay tessellations of the galaxy distribution.  The three reconstruction methods are analysed and compared using mock magnitude-limited and volume-limited SDSS redshift surveys, obtained on the basis of the Millennium simulation. We investigate error trends, biases and the topological structure of the resulting fields, concentrating on the void population identified by the Watershed Void Finder. Environmental effects are addressed by evaluating the density fields on a range of Gaussian filter scales.  Comparison with the void population in the original simulation yields the fraction of false void mergers and false void splits.  In most tests DTFE, NNFE and Kriging have largely similar density and topology error behaviour. Cosmetically, higher order NNFE and Kriging methods produce more visually appealing reconstructions. Quantitatively, however, DTFE performs better, even while computationally far less demanding. A successful recovery of the void population on small scales appears to be difficult, while the void recovery rate improves significantly on scales $> 3\\Mpch$. A study of small scale voids and the void galaxy population should therefore be restricted to the local Universe, out to at most 100$\\Mpch$. ",
        "introduction": "Over the past thirty years a clear paradigm has emerged as large redshift surveys opened the window onto the distribution of matter in our Local Universe: galaxies, intergalactic gas and dark matter exist in a wispy weblike spatial arrangement consisting of dense compact clusters, elongated filaments, and sheetlike walls, amidst large near-empty void regions, with similar patterns existing at higher redshift, albeit over smaller scales.  The Cosmic Web is the fundamental spatial organization of matter on scales of a few up to a hundred Megaparsec, scales at which the Universe still resides in a state of moderate dynamical evolution \\cite{Peebles80,Zeldovich70,Bondweb96}.  Its appearance has been most dramatically illustrated by the recently produced maps of the nearby cosmos, the 2dF galaxy redshift survey (2dFGRS), the Sloan Digital Sky Survey (SDSS) and the 2MASS redshift surveys \\cite{Colless03,Tegmark04,Huchra05}.  According to the standard lore of structure formation, structures emerged from small perturbations in the primordial field of Gaussian density and velocity perturbations. Under the force of gravity these fluctuations grow and cluster to become the present day observed structures. At large scales the density field has been evolving (quasi)-linearly and still retains much information on cosmological parameters and structure formation. The linear density field provides an abundance of probes from which cosmological information can be extracted. Prominent probes are the clustering of galaxies, which has been used to infer the underlying primordial power spectrum of density fluctuations \\citep{Peebles80,Percival01,Tegmark04}, the temperature and polarisation anisotropies in the Cosmic Microwave Background \\citep[e.g.][]{Smoot92,Spergel03} and the shearing of galaxy images by the gravitationally lensed photon paths through the inhomogeneous matter distribution \\citep{Mellier99,Massey07,Hoekstra08}.  While these (quasi-)linear cosmological probes have yielded an impressive amount of cosmological information, the exploitation of the pronounced nonlinear patterns of the cosmic web towards probing cosmological parameters and cosmic structure formation has been less fortuitous. Even though the morphology, shape and other statistical characteristics of the quasi-linear cosmological density field forms a direct reflection of the structure assembly process in the Universe, on the corresponding small scales nonlinear growth has significantly altered and erased some of the essential cosmological information. The absence of an objective and quantitative procedure for identifying and isolating clusters, filaments and voids in the cosmic matter distribution has been a major obstacle in investigating the structure and dynamics of the Cosmic Web. The overwhelming complexity of the individual structures and their connectivity, the huge range of densities and the intrinsic multi-scale nature prevent the use of simple tools that may be sufficient in less demanding problems. However, various interesting new approaches and methods have been forwarded in the past few years, often based on ideas stemming from image processing, mathematical morphology, and medical imaging \\citep[e.g.]{Aragon07,Aragon10,Sousbie10,Way11,Genovese10}.  \\begin{figure*} \\centering \\includegraphics[width=0.9\\textwidth]{sdssdens_test.rev.fig1.eps} \\caption{A visualisation of the (DTFE) SDSS density field (1$\\Mpch$), the contour levels are divided roughly into the overdense and the underdense regime. Both the galaxies (blue dots) and the density represent a slice of 12$\\Mpch$ thickness. Some of the most prominent features have been named, like the {\\it Bo\\\"otes SuperVoid} \\citep{Kirshner81} and the large supervoid ({\\it BS SuperVoid}) identified by \\citet{Bahcall82}. Also the largest overdense structure, the {\\it(Coma) Great Wall}, and the location of the {\\it Hercules supercluster} within the Wall are indicated.} \\label{fig:nicefigure} \\end{figure*}  This study is the first in a series in which we systematically investigate the web-related structures found in the Sloan Digital Sky Survey (SDSS). It involves a systematic program in which we explore the cosmography of the Local Universe, assess the statistical characteristics of the density field, identify and categorize the voids and filaments within the SDSS galaxy redshift sample, and study the biasing of the galaxy population with respect to the mass distribution and the dependence of galaxy properties on the large scale environment.  The intention of this paper is the reconstruction of the underlying (nonlinear) density field by translating the spatially irregularly distributed and discrete galaxy sample in the SDSS galaxy redshift survey into a representative density field.  The density field reconstructions described in this study will form the basis of a series of studies in which we address the statistical and topological properties of the cosmic density field in our Local Universe and in which we analyze the cosmography, void population and spinal structure of the local Cosmic Web in the Sloan Digital Sky Survey. The cosmographic description of the nearby Universe includes the identification and cataloguing of the filaments, voids and clusters in the SDSS galaxy distribution. The key motivation for the density field reconstruction techniques should therefore be that it allows us to probe the complex quasi-linear (and nonlinear) structures that we find in galaxy redshift surveys. This involves the ability to reproduce the distinct anisotropic - filamentary and wall-like - features that make up the Cosmic Web, as well as the ability to trace the hierarchical substructure of the cosmic matter distribution and the ability to identify the void population that forms one of its most salient features.  We are specifically interested in the performance of the fast and efficient linear tessellation-based DTFE density field reconstruction technique \\citep{Schaapwey00,Schaap07,Weyschaap09}. In addition, we investigate two additional higher-order techniques which are potentially suited for representing the quasi-linear density field of the cosmic web, the local natural neighbour interpolation (NNFE) and a nonlocal Kriging interpolation technique, the Natural Lognormal Kriging formalism. By means of an extensive comparison we evaluate which aspects of the density field are best reproduced by either DTFE, NNFE or Natural Lognormal Kriging, and which of these methods is best suited to function for further structural analysis.  In order to be able to assess the reliability of the results of our structural tools, it is crucial to understand the details and errors in reconstructed density fields. We will therefore present a detailed comparison study between three different reconstructions methods, and assess their density and topological errors over a range of scales. This will range form the small nonlinear scales, at ~1$\\Mpch$, to a scale of ~10$\\Mpch$, which represents the transition from quasi-linear to the linear regime.  We focus on the translation of the galaxy positions in the SDSS DR6 galaxy redshift survey to a representative density field within the survey volume. Currently, SDSS encloses the largest and deepest contiguous region of the nearby Universe mapped by a galaxy redshift survey. This makes the SDSS an ideal data sample for a full three-dimensional density field reconstruction. A first impression of the resulting DTFE density field map of a region in the SDSS DR6 survey is shown in figure~\\ref{fig:nicefigure}. Our techniques will be generally applicable to any uniform galaxy redshift survey. Even though survey limits and scales are different from that of the SDSS, it will be straightforward to carry over the results to other redshift surveys.  \\subsection{Inferring Cosmic Density Fields} The richest source of data for investigating the intricate web-like cosmic matter distribution is the distribution of galaxies in galaxy redshift surveys.  We shall make the assumption that the galaxy distribution is a representative tracer of the underlying mass density field and of the underlying structure.  Different samples selected from such catalogues may of course trace different structures.  We set out to explore high fidelity methods for reconstruction of the density field from the discrete galaxy distribution. We require accurate local density estimates that can be used to reliably compute structural and topological indicators. Computational efficiency is an important factor since  our reconstruction technique has to be able to deal with galaxy samples of a million of galaxies as well as with N-body simulations comprising orders of magnitude more particles.  There are two immediate and important implications and complications with respect to our intention of including the (quasi)-nonlinear components in the density field. The first one is that we have to be aware of the noise components that tend to be included in the reconstructions as well \\citep{Schaap07,Weyschaap09}. The second one is that nonlinear data are typically not well behaved, marked by strong gradients in the density field. It means that we have to take special care of those locations marked by non-linearities in the data.  \\subsection{Spatial Point Processes and Continuous Density Fields} We pursue the reconstruction of a density field from a spatial point process consisting of irregularly distributed points.  We assume that the local intensity of the points is a fair tracer of the density and that these values are samples of a continuous underlying density field.  In doing so we appreciate that differently defined samples may trace different structures.  The reconstruction problem we address here is a data processing problem.  The interpretation of the results is an astronomical process.  When the spatial point process is defined by the galaxy distribution, the situation will be more complex. The cosmic web is marked by a distinct luminosity, colour and morphology segregation. A strong and systematic trend of galaxy morphology with density has been established a few decades ago by \\cite{Dressler80}.  Early type galaxies preferentially in rich groups and clusters and late types galaxies residing mainly in filaments and walls \\citep[e.g.][]{Giovanelli86}. Other strong systematic clustering trends with luminosity and colour have also been established, such as in the complete SDSS DR7 galaxy sample \\citep{Zehavi10}. A fair sample of galaxies should in principle take this intrinsic segregation into account. We are pursuing this in a forthcoming publication. In this study, mainly intent on establishing our reconstruction technology, we will consider the total galaxy population. The reconstructed density maps are therefore unlikely to represent a fair reflection of the underlying dark matter matter network, but will nonetheless convey the overall pattern of the large scale structure.  The reconstruction consists of two fundamental steps. The first is the estimation of the local galaxy density. The second is the interpolation of these density values to obtain a continuous spatial density field.  Following straightforward grid based interpolation methods and more sophisticated adaptive filter techniques, there has been a substantial investment into developing more advanced techniques. Examples of recently introduced techniques to follow the multiscale nature of the Cosmic Web is the Multiscale Morphology Filter \\citep{Aragon07} and the Hierarchical SpineWeb technique \\citep{Aragon10}, a Morse theory based formalism that is closely related to the Watershed Void Finder by \\cite{Platen07}. An example of an alternative route involves the wavelet reconstruction by \\cite{Martinez05}.  Here we concentrate in particular on the density estimation and the subsequent field interpolation, in anticipation of the post-processing and feature detection studies of the SDSS DR6/DR7 survey, the subject of the additional papers in this series. To deal with the shot-noise and necessary selection effects in the resulting density fields, the analysis will involve filtering, a necessary post-processing step.  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{sdssdens_test.rev.fig2.ps} \\caption{The schematic outline of our cosmological reconstruction procedure. The {\\it Point Sample} is discussed in section~\\ref{sec:data5}. The {\\it Density Estimation} item will be treated in section~\\ref{sec:density} \\& \\ref{sec:densdtfe} and the {\\it Density Interpolation} methods are introduced in section~\\ref{sec:dtfe},~\\ref{sec:nnfe} and \\ref{sec:kriging}. Section~\\ref{sec:quan}, ~\\ref{sec:topo} and ~\\ref{sec:sdssden} deal with the {\\it Post-Processing}.} \\label{fig:outline} \\end{figure}  \\subsubsection{Density Estimation} \\label{sec:density} The first step in the reconstruction procedure is density estimation.  We show how this fits into our general scheme in Fig.~\\ref{fig:outline}. A similar, but slightly different outline can be found in \\cite{Kitaura08}.  In the literature we may find a large variety of density estimators \\citep[e.g.][]{Silverman86,Sain02,Katkovnik02,Miller03}. Usually they involve filters of some kind. Dependent on the filter kernel, density estimates may have a local character or include the information from a wider range of points in a point distribution. Filter kernels may have a rigid scale and shape, or they may be adapting themselves to the local point density. An example of a global estimator is a Gaussian kernel, which takes along the information - be it weighted - from distant points.  A well-known example of rigid local estimators are the CIC and TSC grid interpolation formalisms \\citep[e.g.][]{Hockney81}. A considerably more flexible and adaptive filter kernel, frequently used in current N-body studies, is that of the spline-based interpolation recipes used in Smooth Particle Hydrodynamics codes \\citep[e.g.][]{Monaghan92} whose scale is determined by the distance to the $K^{th}$ nearest neighbour.  Recently, various local and adaptive density estimators have been shown to provide highly favourable results for complex cosmological matter distributions, marked by a large dynamics range of scales and density values and intricate geometric patterns. One of these exploits the adaptivity of the kD tree \\citep{Ascasibar05,Sharma06}. We also note the use of the Epanechnikov kernel estimator in the identification of superclusters  \\cite{Einasto07}.  For the three reconstruction techniques addressed in this study, we use the local DTFE density estimate introduced by \\cite{Schaapwey00} (see appendix~\\ref{app:dtfe}). It sets the density value at a sample point proportional to the inverse volume of the contiguous Voronoi cell around a sample point, i.e.. the sum of the Delaunay tetrahedra to which the sample point belongs \\citep{Schaapwey00,Schaap07,Weyschaap09}. Tessellation-based methods are based on the realization that the optimal estimate for the spatial density at the location of a point ${\\bf x}_i$ in a discrete point sample ${\\cal P}$ is given by the inverse of the volume of the corresponding Voronoi cell \\cite{Okabe2000}. They have been introduced by \\cite{Brown65} and \\cite{Ord78} and were first used in astronomy, for the specific purpose of devising source detection algorithms, by \\cite{Ebeling93}.  In an extensive study, \\cite{Schaap07} demonstrated that the DTFE estimates are substantially better than those of the rigid grid based CIC and TSC techniques. Also, it outperforms the adaptive SPH spline performance, in particular in areas with substantial density gradients.  \\subsubsection{Interpolation} \\label{sec:interpolation} Interpolation on randomly scattered points aims to approximate a continuous function constrained by the available data points.  A wide variety of approaches have been put forward, each with their own advantages and disadvantages. These methods can be roughly divided into two categories, global and local methods. Local methods have the benefit that they are fast and able to deal with large data-sets. Global methods tend to produce smoother interpolated functions but are computationally more expensive.  An overview and discussion of spatial interpolation techniques can be found in \\cite{Press07} (chapter 3.4) and in \\cite{Watson92}. We refer to \\cite{Lombardi02} for a detailed review of the statistical properties of a large number of techniques, and to \\cite{Franke82} and \\cite{Amidror02} for a detailed comparison between various methods.  Amongst the more sophisticated interpolation techniques we can distinguish various classes: inverse distance based methods (IDW), moving least squares methods, the class of radial basis function interpolation techniques (RBF), Kriging interpolation methods \\citep[see sect.~\\ref{sec:kriging}]{Krige51,Matheron63}) and triangulation based methods, both the linear DTFE technique \\citep[][sect.~\\ref{sec:dtfe}]{Schaapwey00,Weyschaap09} and the higher order NNFE natural neighbour interpolation technique \\citep[][sect.~\\ref{sec:nnfe}]{Sibson81,Watson92}.  \\bigskip RBF and Kriging interpolation methods use predefined kernels to interpolate the field. The kernel in RBF methods is a basis function that spans the space of all the interpolating functions. Kriging interpolation uses spatial correlations between sample points, with its kernel being equal to the corresponding covariance function. The method, introduced by \\cite{Matheron63}, is the {\\it best linear unbiased estimator} of a density value given a set of measured sample points at irregularly spaced points. Given the commonly accepted fact that the primordial density perturbation field is Gaussian, we may therefore see the Kriging interpolator as the natural choice for reconstructing the cosmological density fields which emerged out of these primordial Gaussian circumstances.  There is a distant relationship of Kriging interpolation to Wiener filtering techniques \\citep{Wiener49,Rybicki92}. However, Wiener filtering is based on a different philosophy than Kriging, in that it includes a model for the noise and is evaluated in Fourier space. Also, classical Wiener filtering is predicated on an underlying Gaussian distribution. As a result, it has the serious disadvantage of suppressing or substantially diluting nonlinear structures of interest. More advanced recent developments and applications of Wiener filters to the reconstruction of the density distribution have largely remedied its capacity for reconstructing the density distribution \\citep{Kitaura08,Kitaura09,Kitaura10}.  A rather different approach is advocated in the triangulation-based interpolation techniques.  Both the linear Delaunay Tessellation Field Estimator \\citep[DTFE][]{Schaapwey00,Schaap07,Weyschaap09} and the higher-order Natural Neighbour Field Estimator \\citep{Sibson81,Watson92,Braun95} use the neighbourhood relationships defined by the Voronoi and Delaunay tessellation of the point sample to establish a fully adaptive, irregular and local interpolation grid. DTFE uses the Delaunay triangulation to reconstruct in a self-adaptive, mass conservative and parameter free way the underlying spatial (density) distribution. In combination with the DTFE tessellation-based sample points density estimates (see previous sect.~\\ref{sec:density}), the DTFE interpolation leads to a volume-covering density field which has been shown to recover the hierarchical as well as the anisotropic morphology of the Cosmic Web \\citep{Schaap07,Weyschaap09}.   \\subsection{DTFE, NNFE and Kriging} The DTFE {\\it Delaunay Tessellation Field Estimator} method will be compared to two other techniques having a similar potential for a proper reconstruction of the cosmic web. The NNFE {\\it Natural Neighbour Interpolation} technique shares the local nature of DTFE, but involves higher order interpolations. The third formalism is {\\it Natural Lognormal} Kriging is a version of Kriging interpolation, and thus involves a nonlocal higher-order interpolation methodology.  The crucial step of importance in our investigation is the interpolation step (see diagram Fig.~\\ref{fig:outline}). Using the same initial sample point density estimates, the first step of the reconstruction procedure, we can assess the relative merits of DTFE, NNFE and Kriging interpolators.  \\subsection{Outline of this study} We start by describing the data samples in section~\\ref{sec:data5}, the SDSS survey sample and a mock galaxy sample which mimics the SDSS, obtained from the Millennium simulation. The local DTFE density estimate at the location of each of the sample galaxies is the subject of section~\\ref{sec:density}. In the subsequent section we present and describe each of the three interpolation techniques investigated in this study, DTFE in subsection~\\ref{sec:dtfe}, NNFE in subsection~\\ref{sec:nnfe} and Lognormal Kriging in subsection~\\ref{sec:kriging}. The comparison throughout the paper is based on one specific sample, the density field reconstruction of the SDSS mock galaxy sample. Following a discussion in sect.~\\ref{sec:qual} of the qualitative appearance of the density maps by each of the three methods, we turn to an intensive quantitative error and quality analysis of the density field in sect.~\\ref{sec:quan}.  In sect.~\\ref{sec:topo} this is followed by an investigation of the topological structure of the weblike galaxy distribution in the survey, mainly based on the void population as traced by the Watershed Void Finder. Section~\\ref{sec:sdssden} presents the density field reconstruction of the SDSS data sample and in sect.~\\ref{sec:summary}.we summarize this study.  Following this first paper in a series will be a statistical study of the density  field, focusing in particular on the one-point probability density function. Subsequently, we will present and discuss the cosmography of the reconstructed local Universe. Later studies will analyze the void population in the SDSS density field, and concentrate on the properties of galaxies as function of the large scale environment as characterized by our technique. ",
        "conclusions": "\\label{sec:summary} This study is the first in a series in which we analyze the structure and topology of the Cosmic Web as traced by the Sloan Digital Sky Survey.  In this study we investigate the ability of three reconstruction techniques to analyze and investigate weblike features and geometries in a discrete distribution of objects. The three methods are the Delaunay Tessellation Field Estimator (DTFE), Natural Neighbour Field Estimator (NNFE) and a local ``natural'' implementation of Kriging, the {\\it Natural Lognormal Kriging} technique. DTFE and NNFE are based on the local geometry defined by the Voronoi and Delaunay tessellations of the galaxy distribution. The Kriging formalism is adapted and optimized for the the approximate lognormal density distribution encountered in the mildly nonlinear cosmic web and is based on the logarithm of the measured density values. Also, we have chosen to restrict the evaluations to a localized neighbourhood, the {\\it tetrahedral natural neighbourhood} based on the Delaunay triangulation of the point sample.  The three reconstruction methods are analysed and compared using mock magnitude-limited and volume-limited SDSS redshift surveys, obtained on the basis of the Millennium simulation. We investigate error trends, biases and the topological structure of the resulting fields. The reconstructed density fields, mainly from the magnitude-limited mock survey samples but also from volume-limited ones, are compared with the density field of the total simulation galaxy sample. The differences between the various field reconstructions are investigated on the basis of an error analysis, mostly involving point-to-point comparisons. Environmental effects are addressed by evaluating the density fields on a range of Gaussian filter scales. With respect to the topology of the survey fields, we concentrate on the void population identified by the Watershed Void Finder \\citep{Platen07}.  The void population in the full density field and in the magnitude-limited survey density fields are compared. The number of {\\it false mergers} - in which original voids emerge as a part of a larger void in the survey field - and of {\\it false splits} - in which an original void splits up in one or more voids in the survey field - forms the basis of the topological quality evaluation.  By investigating the quality of the resulting density field estimates, over a range of scales and in different environments, as well as the more global topological structure of the weblike network, we wish to identify and understand the qualities of these techniques for the different purposes and post-processing steps which are the subject of the following papers in this series. The following observations were made:  \\begin{itemize} \\item[$\\bullet$] In most tests, DTFE, NNFE and Kriging have largely similar density and topology error behaviour. \\item[$\\bullet$] Cosmetically, higher order NNFE and Kriging methods produce more visually appealing reconstructions. \\item[$\\bullet$] Quantitatively, DTFE performs (marginally) better. Part of this at first sight surprising finding is a consequence of the higher sensitivity of the higher order NNFE and Kriging interpolation to intrinsic errors in the galaxy sample. An additional factor is the smaller natural neighbourhood of DTFE and NNFE with respect to the 3 to 4 times larger neighbourhood of Natural Lognormal Kriging, which restricts density errors to smaller volumes. \\item[$\\bullet$]  With respect to the topological properties of the reconstructed density fields, it has become clear that a successful recovery of the void population on small scales is rather difficult. On these scales, the removal of only a couple of void galaxies leads to the spurious merging of observed voids. \\item[$\\bullet$] The void recovery rate improves significantly at filter scales $> 3\\Mpch$. The immediate repercussion is that a proper analysis of small scale voids, and void galaxies, has to be necessarily restricted to the local Universe out to at most 100$\\Mpch$. As the environmental influences on the galaxy formation process seem to be mostly determined on these scales \\citep{Park07b}, our study within this project, subject of a forthcoming paper in this series, will restrict itself mainly to this volume. \\end{itemize}  \\bigskip A variety of technical improvements of our DTFE, NNFE and Kriging implementations may lead to a better performance. One immediate option is to invoke an image grid which has a more natural character for the galaxy survey context of our study than the cubic grid we have used in the evaluations described in this paper. Because of the flexibility of their definition, the Kriging formalism will be form a promising context for further adaptations and improvements. Of immediate importance is the implementation of non-local techniques to optimize the matrix inversion calculations while retaining the influence of large scale correlations and predictor-corrector methods to deal with oscillatory instabilities. While we have not yet investigated its performance, results of studies in other fields emphasize the importance of evaluating the performance of Radial Basis Function techniques as a possible alternative.  \\bigskip The DTFE, NNFE and Kriging density field reconstructions form the basis of a series of studies in which we analyze the cosmography, void population and spinal structure of the local Cosmic Web in the Sloan Digital Sky Survey. Given the ease and efficiency of calculation, and its good quantitative behaviour, for most of these studies we use the DTFE formalism. However, the Natural Lognormal Kriging results looks very promising and appears to produces a well-behave coherent weblike density map of the SDSS survey. With the large advantage of controlling the error behaviour and properties of the reconstructed map, along with the large potential for extensions and optimizations of the method, the Natural Lognormal Kriging maps will play a dominant role in our study of the Cosmic Web.  \\bigskip The density field reconstructions within the SDSS DR6, and subsequently SDSS DR7, volumes will be subjected to a statistical study in the second paper of this series. We will focus in particular on the 1-point probability function of the SDSS density field. With the help of corresponding magnitude-limited mock catalogues we will infer in how far the lognormal probability function, as well as related higher order versions, form a proper description of its statistical character. Also, the mock catalogues will allow us to assess the bias of the galaxy population in high density areas and, in particular, the low density void region. A cosmographic description of the reconstructed Local Universe, along with the identification of filamentary supercluster complexes, voids and supervoid complexes is the subject of the third paper in this series. The study of the void population in the SDSS density field, concerning a detailed assessment of the void size and shape properties, forms an important rationale behind the development of the density field reconstructions described in this study. This has become an interesting area of research, in particular as recent studies have emphasized the potential of extracting cosmological information from cosmic voids, in particular that concerning the dark energy equation of state \\citep{LeePark09,Lavaux09,Biswas10}. In recent years, there has also been a strong interest in large scale environmental influences on galaxy properties and on the galaxy formation process. The tools described in the present study, will allow an assessment on the basis of a properly defined density field on quasi-linear scales."
    },
    "1107/1107.1993_arXiv.txt": {
        "abstract": "{} {We present a study of the temporal evolution of coronal loops in active regions and its implications for the dynamics in coronal loops.} {We analyzed images of the Atmospheric Imaging Assembly (AIA) on the Solar Dynamics Observatory (SDO) at multiple temperatures to detect apparent motions in the coronal loops.} {Quasi-periodic brightness fluctuations propagate upwards from the loop footpoint in hot emission at 1~MK, while sporadic downflows are seen in cool emission below 1~MK. The upward motion in hot emission increases just after the cool downflows. } {The apparent propagating pattern suggests a hot upflow from the loop footpoints, and is considered to supply hot plasma into the coronal loop, but a wavelike phenomenon cannot be ruled out. Coronal condensation occasionally happens in the coronal loop, and the cool material flows down to the footpoint. Emission from cool plasma could have a significant contribution to hot AIA channels in the event of coronal condensation.} ",
        "introduction": "Heating of a coronal loop is a fundamental problem in solar physics. Coronal loops, which are thin thread-like structures seen in coronal emission lines, must be continuously heated to account for their radiative and conductive energy losses. However, models of coronal loops are not mature enough to reproduce the observations sufficiently \\citet{winebarger2002},\\citet{warren2006}, and \\citet{aschwanden2009} show that a static loop model is not consistent with observations. A dynamic model with a collection of impulsive heatings is in better agreement with observations, but still not satisfactory \\citep{warren2003,warren2007}. A detailed study of the temporal evolution of coronal loops is therefore crucial for understanding them.  Observations of nonflaring active regions show that coronal loops are visible over a wide temperature range and are highly variable. Brightness distribution of EUV emission lines indicates that hot loops ($\\geq 2\\times 10^6$~K) are concentrated in the cores of active regions, while lower temperature loops ($\\sim 1\\times 10^6$~K) are located on the periphery of the active region \\citep{aschwanden2008c,tripathi2008,odwyer2011}. Loop-like structures are also seen at transition-region and chromospheric temperatures. \\citet{kjeldsethmoe1998} have shown that cool loops seen at $1\\sim5 \\times 10^5$~K change significantly within one hour, while hot loops are less variable. \\citet{schrijver2001} finds cool plasma ($\\leq1\\times 10^5$~K) flowing down along the loop, which he interprets as a result of catastrophic cooling at the loop top. Numerical simulations shows that frequent occurrence of falling blobs in chromospheric lines, or coronal rain, favors a heating concentrated near the footpoints of a coronal loop \\citep{mueller2003,mueller2004,mueller2005,antolin2010}. Clearly, in order to understand the energy balance of coronal loops, it is important to study both the cooling and heating processes.  The Atmospheric Imaging Assembly \\citep[AIA;][]{lemen2011,boerner2011} on the Solar Dynamics Observatory (SDO) records high-cadence EUV images at multiple temperatures. In this paper, we selected nonflaring active regions at the limb and on the disk to study the temporal evolution of coronal loops. Unprecedented data from AIA allowed studying the flows along the loops over a wide temperature range.  The paper is organized as follows. Observations with AIA are described in Sect.~2. Section 3 shows the temporal evolution of selected loops in the active region and simulated lightcurves of cooling plasma under constant pressure and Sect~4 summarizes the results and discusses the implications for coronal loop models.    \\begin{figure*} \\centering \\includegraphics[width=16cm]{fig1.eps} \\caption{Active region NOAA 11117 at the limb as seen in AIA-131\\AA\\, and 193\\AA\\, channels on 31 October 2010. The two right panels show temporal variations between 13:15 UTC and 13:25 UTC (See Sect. 3.1). The cross signs {\\it A} and {\\it B} indicate darkening in AIA-193\\AA\\, coinciding with brightenings in AIA-131\\AA.} \\label{fig_limb} \\end{figure*}  \\begin{SCfigure*} \\includegraphics[width=14cm]{fig2.eps} \\caption{The active region NOAA 11120 on the disk {\\it Left:} AIA-193\\AA\\, image. A dashed line indicate the location of time distance plots (Fig. \\ref{fig_slice_m}). {\\it Middle:} Doppler map deduced from \\ion{Fe}{xii}~$\\lambda$195.12~\\AA\\, spectrum. {\\it Right:} Doppler map in \\ion{Si}{vii}~$\\lambda$275.35~\\AA.} \\label{fig_diverge} \\end{SCfigure*} ",
        "conclusions": "AIA data in active regions reveal brightness variations propagating both upward and downward. In AIA-193\\AA, the variations are always present and propagate upward with the speeds of 57~km~s$^{-1}$ to 88~km~s$^{-1}$. Our study from different viewing angles indicates that these fluctuations originate at the footpoints of the loop system and propagate along magnetic field lines diverging from the footpoint. The blueshift observed at loop footpoints supports an interpretation in terms of a constant hot upflow, although the magnitude of the Doppler shift (up to $-20$~km~s$^{-1}$) is smaller than the propagation speed along the investigated cut in AIA-193\\AA. The low velocity could result from a superposition of different velocity components under the spatial resolution of EIS. It has been reported in hot emission lines (2~MK) that the blue wing of the line profile is enhanced in active regions, which is interpreted as a signature of subsonic upflows \\citep{hara2008a,depontieu2009,peter2010}. A detailed study of spectral profiles at the coronal loop footpoints will be presented in a future paper. We cannot rule out that the quasi-periodic intensity fluctuation is caused by a propagating magneto-acoustic wave (Fig.~\\ref{fig_slice_m}); however, \\citet{depontieu2010} demonstrated that lightcurves from randomly driven upflows and propagating waves are indistinguishable.  Sporadic downflows towards the loop footpoint are detected in the cooler AIA-304\\AA, 131\\AA, and 171\\AA\\, channels, which we interpret as a cool plasma flowing down along the magnetic field lines. The downflow of cool emission is reported in previous studies \\citep{schrijver2001, ugarte-urra2009}. Our study revealed that the downflowing plasma is initiated by a cooling in the corona. The increase in AIA-171\\AA\\, and AIA-131\\AA\\, lightcurves coincide with a small decrease in the AIA-193\\AA\\, lightcurve (Fig. \\ref{fig_lc2}). The results support the idea of coronal condensation occurring in the coronal loop \\citep{mueller2003,mueller2004,mueller2005}. However, AIA-193\\AA\\, increased again and remained high, while AIA-131\\AA\\, decreased. Although the observed sequence of peak emissions in multiple AIA channels does not agree with the order of maximum temperature response of the AIA channels, it can be qualitatively explained by a cooling plasma at a constant pressure (Fig. \\ref{model}). The simulated lightcurve of AIA-193\\AA\\, shows that the contribution of cool emission lines within the broad passbands of the AIA channel is not negligible in the event of coronal condensation. Care should be taken when interpreting the lightcurves of AIA. There is also a need for more realistic modeling of cooling loops, which can be used to carry out a more quantitative comparison with our observations.  The descent of cool plasma is seen in AIA-304\\AA\\, first, and is followed by AIA-171\\AA\\, (Fig.\\ref{fig_slice}). It seems a cool dense core in AIA-304\\AA\\, is surrounded by diffuse hotter plasma in AIA-171\\AA. This might be a reason for long decay times of lightcurves from AIA-131\\AA, 171\\AA, and 193\\AA. It is interesting to note that the propagation pattern in AIA-193\\AA\\, becomes more pronounced after a falling event in cool emission (Figs. \\ref{fig_slice_m} and \\ref{fig_slice}). A possible explanation is that the coronal loop is emptied after the descent of cool plasma, so there is plenty of room for hot upflows to propagate into the coronal loop. This work demonstrates that a dynamic modeling is essential for understanding the nature of coronal loops."
    },
    "1107/1107.4608_arXiv.txt": {
        "abstract": "\\noindent We have searched $\\approx$8200 sq.~degs for high proper motion ($\\approx$0.5--2.7\\arcsec/year) T~dwarfs by combining first-epoch data from the Pan-STARRS1 (PS1) \\TPS, the 2MASS All-Sky Point Source Catalog, and the \\WISE\\ Preliminary Data Release. We identified two high proper motion objects with the very red $(W1-W2)$ colors characteristic of T~dwarfs, one being the known T7.5~dwarf GJ~570D. Near-IR spectroscopy of the other object (\\psobjectfull~$\\equiv$~\\wiseobjectfull) reveals a spectral type of T8, leading to a photometric distance of $7.2\\pm0.7$~pc. The 2.56\\arcsec/yr proper motion of \\psobject\\ is the second highest among field T~dwarfs, corresponding to an tangential velocity of $87\\pm8$~\\kms. According to the Besan\\c{c}on galaxy model, this velocity indicates its galactic membership is probably in the thin disk, with the thick disk an unlikely possibility. Such membership is in accord with the near-IR spectrum, which points to a surface gravity (age) and metallicity typical of the field population. We combine 2MASS, SDSS, \\WISE, and PS1 astrometry to derive a preliminary parallax of $171\\pm45$~mas ($5.8^{+2.0}_{-1.2}$~pc), the first such measurement using PS1 data. The proximity and brightness of \\psobject\\ will facilitate future characterization of its atmosphere, variability, multiplicity, distance, and kinematics. The modest number of candidates from our search suggests that the immediate ($\\sim$10~pc) solar neighborhood does not contain a large reservoir of undiscovered T~dwarfs earlier than about T8. ",
        "introduction": "The \\ps\\ (PS1) survey \\citep{PS1_system} represents the first of a new generation of multi-epoch digital sky surveys. In an effort to identify nearby T~dwarfs and rare ultracool dwarfs of interest, we are conducting a proper motion search that combines the first epoch of PS1 data with 2MASS, with initial results from PS1 commissioning data in \\citet{deacon11-ps1-tdwarfs}. In this Letter, we present the discovery of a very high proper motion late-T~dwarf, \\psobjectfull\\ (hereinafter \\psobject). This object has independently been identified by \\citet{2011A&A...532L...5S} in a search for bright, very red objects in the \\WISE\\ Preliminary Data Release with proper motions of $\\gtrsim$0.3\\arcsec/year. They estimated a spectral type of T8--T10 based on the $(W1-W2)$ color, a photometric distance of 5.5$^{+2.3}_{-1.6}$~pc, and possible thick disk membership. Our near-IR spectroscopy gives a spectral type of T8, a photometric distance of $7.2\\pm0.7$~pc, and probable thin disk membership. ",
        "conclusions": "For the median \\vtan\\ of 30~\\kms\\ of ultracool dwarfs within 20~pc \\citep{2009AJ....137....1F}, our proper motion range of 0.5--2.7\\arcsec/yr corresponds to distances of 12~pc and 2~pc. (Objects with higher tangential velocities of course can be selected at even larger distances.) This approximate kinematic limit is well-matched to the $J\\approx16.5$~mag limit for our search: the three T7.5~dwarfs with 2MASS photometry and parallaxes have $M_J$~=~15.93, 15.65, and 16.47~mag (HD~3651B, 2MASS~J1217$-$03, and GJ~570D, respectively) and the T8~dwarf 2MASS~J0415$-$09 has $M_J=16.90$~mag, meaning our search has a limiting distance of $\\approx$12~pc for T8~dwarfs (9~pc for the 2MASS full-sky completeness of $J=15.8$). The limiting distance drops precipitously for later types: the T10~dwarf UGPS~J0722$-$05 \\citep{2010MNRAS.408L..56L} at 4~pc has $J=16.49\\pm0.13$~mag, right at the nominal 2MASS limit.  The modest number of candidates from our search suggests that the immediate ($\\sim$10~pc) solar neighborhood does not contain a large reservoir of undiscovered T~dwarfs down to about T8. We found 7--10~candidates (depending on the cutoff used for the PS1-\\WISE\\ separation) with proper motions of 0.5--2.7\\arcsec/year and colors of $(W1-W2)>0.7$, including 3~previously known objects. In comparison, Dwarfarchives.org lists 30 objects in the same proper motion range and with spectral types of T0--T8 (ignoring 2 tight companions found by adaptive optics imaging). Since our PS1+2MASS+\\WISE\\ search covered $\\approx$20\\% of the sky, a rough estimate based on the known census predicts we should identify about 6~objects, consistent with our findings."
    },
    "1107/1107.4043.txt": {
        "abstract": "The exact solution for the shape and gravitational field of a rotating two-layer Maclaurin ellipsoid of revolution is compared with predictions of the theory of figures up to third order in the small rotational parameter of the theory of figures. An explicit formula is derived for the external gravitational coefficient $J_2$  of the exact solution. A new approach to the evaluation of the theory of figures based on numerical integration of ordinary differential equations is presented. The classical Radau-Darwin formula is found not to be valid for the rotational parameter $\\epsilon_2 = \\Omega^2/(2\\pi \\mathrm{G}\\rho_2) \\geq 0.17$ since the formula then predicts a surface eccentricity that is smaller than the eccentricity of the core-envelope boundary. Interface eccentricity must be smaller than surface eccentricity. In the formula for $\\epsilon_2$, $\\Omega$ is the angular velocity of the two-layer body, $\\rho_2$ is the density of the outer layer, and G is the gravitational constant.  For an envelope density of 3000~kg~m$^{-3}$ the failure of the Radau-Darwin formula corresponds to a rotation period of about 3~hr. Application of the exact solution and the theory of figures is made to models of Earth, Mars, Uranus, and Neptune. The two-layer model with constant densities in the layers can provide realistic approximations to terrestrial planets and icy outer planet satellites. The two-layer model needs to be generalized to allow for a continuous envelope (outer layer) radial density profile in order to realistically model a gas or ice giant planet. ",
        "introduction": "\\label{sec:intro} \\citet{KongZhangSchub:JGR2010} have presented an exact theory for the rotational distortion of a rotating two-layer spherical body with a constant density core surrounded by an envelope (outer layer) with a different constant density. The solution for the case when the core and envelope have equal densities, the Maclaurin ellipsoids, was obtained more than 250 years ago and attracted the attention of such notables as d'Alembert,  Clairaut, Euler, Laplace, Legendre, Poisson and Gauss. The solutions are discussed in \\citeauthor{Chand:EFE1969} (\\citeyear{Chand:EFE1969}, Ellipsoidal Figures of Equilibrium) and \\citeauthor{Lamb:H1932}~(\\citeyear{Lamb:H1932}, Hydrodynamics). Until the work of \\citet{KongZhangSchub:JGR2010}, the classical solution for the constant density Maclaurin ellipsoid had not been generalized to a body with non-uniform density. Instead, approximate solutions, for bodies with density increasing from the surface to the center, have been developed by geophysicists interested in the internal structures of the Earth and planets. The approximate solutions fall under the umbrella of the theory of figures \\citep[][Physics of Planetary Interiors]{ZharkTrubi:PPI1978} and rely on the smallness of a rotational parameter that measures the distortion of a rotating body. In this paper we refer to the theory developed by \\citet{ZharkTrubi:PPI1978} for the shapes of fluid bodies as the theory of figures. The rotational distortions of the Earth and planets are indeed small, but it is important to accurately determine the small distortions to correctly infer the interior structures of the bodies. Accordingly, the theory of figures has been developed to high order in the small rotational parameter.\\par  While the two-layer spherical body with constant core and envelope densities is too simple a model for the careful study of most planets, the exact solution for the rotational distortion of such a body provides a benchmark against which the accuracy and range of validity of approximate solutions and numerical models can be evaluated. In this paper we compare results from the exact solution for the rotating two-layer Maclaurin ellipsoids to those obtained using the Radau-Darwin approximation and other simplifications based on the theory of figures. We derive formulae for the rotationally distorted two-layer sphere using the theory of figures valid to third order in the small rotational parameter and assess these results against the exact solution. Application of the exact theory is made to the Earth and planets while keeping in mind the limitations of using a two-layer model to represent these bodies. Both the theory of figures and the exact theory of \\citet{KongZhangSchub:JGR2010} assume hydrostatic equilibrium of the rotating body.\\par ",
        "conclusions": "The exact solution for the rotational distortion of a two-layer Maclaurin ellipsoid reported in \\citet{KongZhangSchub:JGR2010} has been extended here to provide formulas for the standard spherical harmonic expansion of the external gravitational field of the body. We have also presented a new approach to the evaluation of the theory of figures based on numerical integration of ordinary differential equations.  The classical Radau-Darwin formula is a low order result from the theory of figures and its realm of validity has been evaluated for the two-layer model using the exact solution. It was found that the Radau-Darwin approximation is not valid for the rotational parameter $\\epsilon_2=  \\Omega^2/(2\\pi \\mathrm{G}\\rho_2)\\geq 0.17$  since the formula predicts a surface eccentricity that is smaller than the eccentricity of the core-envelope boundary. Interface eccentricity must be smaller than surface eccentricity. For an envelope density of 3000~kg~m$^{-3}$ the failure of the Radau-Darwin formula corresponds to a rotation period of about 3~hr.  The generalized Roche model, a two-layer model with an envelope density equal to zero, provides a simple model against which to evaluate the validity of the theory of figures against the exact solution. It was found that the theory of figures only slightly underestimates the eccentricities of the surface and core-envelope interface compared with the exact solution.  Application of the exact solution and the theory of figures is made to models of Earth, Mars, Uranus, and Neptune. It is found that the two-layer model with constant densities in the layers can provide realistic approximations to terrestrial planets and icy outer planet satellites. This is perhaps not surprising since the zeroth order structure of these planetary bodies is similar to the two-layer model with constant densities in the layers. The situation is not as straightforward for giant planets since a constant density envelope is not a particularly good representation of the density in the outer layers of such planets. However, the theory of figures, as developed in this paper, is readily generalized to models with arbitrary radial density profiles in the envelope (though we have not carried this out in this paper). Such models will be particularly useful for Jupiter and Saturn which might possess heavy element cores surrounded by gaseous envelopes. The envelope density can be represented by polynomial functions of radius. Inversions of gravitational data based on these models provide constraints on the gas giant interiors independent of assumptions about composition and equations of state. The exact solution for the two-layer Maclaurin ellipsoid can also be extended to allow for a non-constant radial profile of envelope density. This is not as straightforward as the generalization of the theory of figures, but it can be done. The solutions for two-layer bodies can therefore provide acceptable models for the rotational distortion of terrestrial, gas giant, and ice giant planetary bodies.  These solutions can also serve as benchmarks to test the validity of complicated numerical models that invert gravitational and shape data to infer the interior structure of planets. \\clearpage %"
    },
    "1107/1107.0231.txt": {
        "abstract": "The aim of this review is to describe the nature, formation and evolution of the three kinds of high mass X-ray binary (HMXB) population: {\\it i.} systems hosting Be stars (BeHMXBs), {\\it ii.} systems accreting the stellar wind of supergiant stars (sgHMXBs), and {\\it iii.} supergiant stars overflowing their Roche lobe. There are now many new observations, from the high-energy side (mainly from the {\\it INTEGRAL} satellite), complemented by multi-wavelength observations (mainly in the optical, near and mid-infrared from ESO facilities), showing that a new population of supergiant HMXBs has been recently revealed. New observations also suggest the existence of evolutionary links between Be and stellar wind accreting supergiant X-ray binaries. I describe here the observational facts about the different categories of HMXBs, discuss the different models of accretion in these sources (e.g. transitory accretion disc versus clumpy winds), show the evidences of a link between different kinds of HMXBs, and finally compare observations with population synthesis models. ",
        "introduction": "Nearly 50 years after the discovery of the first extra-solar X-ray source \\citep[Sco\\,X-1;][]{giacconi:1962}, X-ray astronomy has now reached its age of reason, with a plethora of telescopes and satellites covering the whole electromagnetic spectrum, obtaining precious observations of these powerful celestial objects.  %\\subsection{Different types of X-ray binaries}  High energy binary systems are composed of a compact object -- neutron star (NS) or black hole (BH) -- orbiting a companion star, and accreting matter from it \\citep[cf e.g.][for a review]{chaty:2006a}. The companion star is either a low mass star (typically $\\sim 1 \\Msol$ or less, with a spectral type later than B, called in the following LMXB for ``Low-Mass X-ray Binary''), or a luminous early spectral type OB high mass companion star (typically $> 10 \\Msol$, called HMXB for ``High-Mass X-ray Binary''). 300 high energy binary systems are known in our Galaxy: 187 LMXBs and 114 HMXBs \\citep[respectively 62\\% and 38\\% of the total;][]{liu:2006,liu:2007}.  Accretion of matter is different for both types of sources\\footnote{I will not review here the so-called $\\gamma$-ray binaries, where most of the energy is in the GeV-TeV range, coming from interaction between relativistic electrons of the pulsar and the wind of the massive star.}. In the case of LMXBs, the small and low mass companion star fills and overflows its Roche lobe, therefore accretion of matter always occurs through the formation of an accretion disc. The compact object can be either a NS or a BH, Sco X-1 falling in the former category. % In the case of HMXBs, accretion can also occur through an accretion disc, for systems in which the companion star overflows its Roche lobe; however this is generally not the case, and there are two alternatives. The first one concerns stars with a circumstellar disc, and here it is when the compact object -- on a wide and eccentric orbit -- crosses this disc, that accretion periodically occurs (case of HMXBs containing a main sequence early spectral type Be III/IV/V star, rapidly rotating, called in the following BeHMXBs). The second case is when the massive star ejects a slow and dense wind radially outflowing from the equator, and the compact object directly accretes the stellar wind through e.g. Bondy-Hoyle-Littleton processes (case of HMXBs containing a supergiant I/II star, called sgHMXBs). %(cf Figure \\ref{lmxb-hmxb}: LMXB on the left and HMXB on the right). We point out that there also exists binary systems for which the companion star possesses an intermediate mass (typically between 1 and $10 \\Msol$, called IMXB for \"Intermediate-Mass X-ray Binary\"). ",
        "conclusions": "Let us first recall the {\\it INTEGRAL} legacy:  \\begin{itemize}  \\item The {\\it INTEGRAL} satellite has quadrupled the total number of known Galactic sgHMXBs, constituted of a NS orbiting a supergiant star. Most of the new sources are slow and absorbed X-ray pulsars, exhibiting a large $\\nh$ and long $P_\\mathrm{spin}$ ($\\sim1$\\,ks).  \\item The {\\it INTEGRAL} satellite has revealed the existence in our Galaxy of two previously hidden populations of high-energy binary systems. First, the SFXTs, exhibiting brief and intense X-ray flares -- with a peak flux of 1 Crab during 1--100s every $\\sim 100$\\,days --, which can be explained by accretion through clumpy winds. Secondly, a previously hidden population of obscured and persistent sgHMXBs, composed of supergiant companion stars exhibiting a strong intrinsic absorption and long $P_\\mathrm{spin}$, with the NS deeply embedded in the dense stellar wind, forming a dust cocoon enshrouding the whole binary system.  \\end{itemize}  Apart from these observational facts, has {\\it INTEGRAL} allowed us to better understand sgHMXBs and other populations of HMXBs? Do we better apprehend the accretion processes in HMXBs in general, and in sgHMXBs in particular, and what makes the fast transient flares so special, in the context of the clumpy wind model, and of the formation of transient accretion discs? In summary, has the increased population of supergiant HMXBs allowed us a better knowledge of these sources, compared to the ones that were already known before the launch of {\\it INTEGRAL}? The answer to all these questions is probably {\\it ``not yet''}, however we now have in hand more sources, and therefore more constraints to play with. Studying these populations will provide a better understanding of the formation and evolution of short-living HMXBs, and study accretion processes. 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. %Our final word is that only a multiwavelength study reveal the nature of these obscured high-energy sources.  %The {\\it Fermi} satellite (Gamma-ray Large %Area Space Telescope, operating in the 10 keV - 300 GeV energy range) %will certainly discover new and unexpected objects.  Indeed, %while these sgHMXBs do not seem to be GeV Emitters, they are %peculiar members of this family which emit at these energy ranges, for %instance LSI +61 303 (a \"disguised\" pulsar, \\citeauthor{dubus:2006a} %\\citeyear{dubus:2006a}), LS 5039 %(\\citeauthor{paredes:2000} \\citeyear{paredes:2000}; %\\citeauthor{aharonian:2005} \\citeyear{aharonian:2005}), %and Cygnus X-1 \\citep{albert:2007}.  This high energy satellite %has a huge potential: 200 individual sources have been detected by %EGRET... we expect that thousands will be detected by Fermi"
    },
    "1107/1107.2923_arXiv.txt": {
        "abstract": "Observations of the Galactic centre (GC) have accumulated a multitude of ``forensic'' evidence indicating that several million years ago the centre of the Milky Way galaxy was teeming with star formation and accretion-powered activity -- this paints a rather different picture from the GC as we understand it today. We examine a possibility that this epoch of activity could have been triggered by the infall of a satellite galaxy into the Milky-Way which began at the redshift of $z=\\newredshift$ and ended few million years ago with a merger of the Galactic supermassive black hole with an intermediate mass black hole brought in by the inspiralling satellite. ",
        "introduction": "There is mounting observational evidence that the epoch that ended several million years ago was marked by an unusual level of activity in the Galactic centre (GC). This is remarkable given that at the current epoch, the GC is best characterized by the quiescent and underluminous nature of ${\\rm Sgr~A^*}$\\citep{genzel10}. The picture of the GC as a once powerful nucleus has begun to emerge from circumstantial observational evidence, most recently strengthened by a discovery of the ``Fermi bubbles'', a pair of giant gamma-ray emitting bubbles that extend nearly 10~kpc north and south of the GC~\\citep{dobler10,su10}. Although there are alternative steady state models for forming the bubbles \\citep{crocker11}, the well defined shock fronts at their edges suggest an abrupt origin. Current explanations include a past accretion event onto the supermassive black hole~\\citep[SMBH;][]{su10,zubovas11}, AGN jets~\\citep{guo11}, a nuclear starburst~\\citep{su10}, and a sequence of star capture events in the last $\\sim$10~Myr~\\citep{cheng11}.  The period of increased gamma-ray activity is consistent with the finding that until several hundred years ago ${\\rm Sgr~A^*}$ was orders of magnitude more X-ray luminous than it is today, as indicated by the echo in the fluorescent Fe~K line emission detected in the direction of the molecular clouds in the vicinity of ${\\rm Sgr~A^*}$ \\citep{inui09,ponti10,terrier10}. Although we cannot be certain that the current quiescence of the GC is unusual, it appears clear that the GC experienced an active phase as recently as a few hundred years ago.  The GC is also a hotbed of star formation containing the three most massive young star clusters in the Galaxy: the Central cluster, the Arches cluster, and the Quintuplet cluster \\citep[see][for a review]{figer08}. The three clusters are similar in many respects. Each cluster contains $\\sim 10^4 \\Msun$ in stars and has central stellar mass density that exceeds those measured in most globular clusters. While our current understanding of massive star and star cluster formation is incomplete, it is plausible that these clusters are all characterized by the star formation event within the past 2$-$7~Myr that resulted in the formation of more massive stars (above $100 \\Msun$) than anywhere else in the Galaxy \\citep{krabbe95,paumard06}. It is possible that the three clusters have a common origin and that they have formed as a consequence of a single event that triggered the flow of the copious amounts of gas into the central $\\sim 50$~pc in the Galaxy \\citep[though see][for a scenario in which the Arches cluster forms at the intersection of X1/X2 gas orbits in the inner Galaxy]{StolteEtAl08}.  On even smaller scales, the existence of massive and young stars in the Central cluster, well within the central parsec, is especially puzzling given their close proximity to the central SMBH. Among those scenarios proposed are: in-situ formation \\citep{bonnell08,mapelli12}, inspiral and consequent disruption of a dense stellar cluster with a central intermediate mass black hole (IMBH) \\citep{merritt09}, and binary disruption by massive perturbers \\citep{perets10}. A clue in favor of the in-situ formation is that most O and Wolf-Rayet type stars at the galactic center seem to inhabit one or more disc-like structures, pointing to their birth in a dense accretion disc~\\citep{bartko10,bonnell08,mapelli12}. Star formation in a gaseous disc also provides a natural explanation for the cuspy distribution of the young stars. In order for starforming clumps to withstand tidal forces in the inner parsec of the GC their densities need to be in excess of $10^{11}\\,{\\rm cm^{-3}}$, at least five orders of magnitude higher than the average density of molecular clouds in the GC~\\citep{figer00}. Such densities can only be achieved through highly compressive events~\\citep{figer08} and it is plausible that both the inflow of large amounts of gas into the GC and its shocking and compression have been caused by a common culprit. Both phenomena are found to arise as consequences of galactic mergers~\\citep{noguchi88,bh91,bh92,bh96,mh96,hq10} making this a possibility worth examining.  Further evidence that the MW has recently survived a dramatic event comes from the distribution of late-type stars in the GC. While the early-type stellar distribution appears to be cuspy~\\citep{genzel03,paumard06,buchholz09,do09,bartko10}, there seems to be a distinct lack of late-type stars. This evidence is based on number counts of spectroscopically identified late-type stars brighter than magnitude $K=15.5$ within the sub-parsec region about Sgr~A$^*$. The best fits of the inner density profile for the late-type stellar population seem to favor power-laws with slopes of $\\gamma < 1$ and even allow the possibility of a core with $\\gamma < 0$, with the stellar density decreasing toward the centre~\\citep{buchholz09,do09,bartko10}. At this stage, the evidence for a deficit of late-type stars is compelling, however there are still significant uncertainties in the density profile: the population of stars on which this inference has been made are luminous late-type giants that comprise only a small fraction of the underlying stellar density of the late-type population. Regardless of the precise slope, however, the distribution of the late-type population is contrasted by the steeply rising density distribution of early-type stars.  Possible mechanisms that could create a core in the distribution of late-type stars have been discussed by~\\citet{merritt10} and include 1) stellar collisions that strip red giants of their envelopes such that they are under-luminous, 2) destruction of stars on orbits that pass close to the SMBH in a triaxial nucleus, 3) inhibited star formation near the SMBH at the time when the late-type population was formed, and 4) ejection of stars by a massive black hole binary. In light of the other evidence that points to a discrete event in the recent history of the GC, it is interesting to revisit the latter mechanism.  In giant elliptical galaxies, the existence of cores is often attributed to ejection of stars by an inspiralling binary SMBH~\\citep{merritt01,faber97,ferrarese06,mm01} and possibly due to gravitational wave recoil after binary coalescence~\\citep{bk04,gm08}. The prediction of these models is that the central stellar mass deficit (traced by the stellar light) is proportional to the mass of the central black hole, $M_{\\rm def}\\propto M_{\\bullet}$ \\citep{graham04,hh10}. This correlation is interesting in view of the observed dichotomy between ellipticals with cores and those with the extra central light: core light deficit was found to correlate closely with $M_{\\bullet}$ and stellar velocity dispersion $\\sigma$, in agreement with the theoretical predictions~\\citep{mm01,gm08}, however, the extra light does not~\\citep{kb09}. An explanation of these phenomena offered by~\\citet{kb09} is that the extra light ellipticals were made in wet mergers with starbursts, where stars formed from gas leftover after the merger, while core ellipticals were created in dry mergers. In galaxies with excess light, the newly formed population of stars fills the core left in the distribution of the older population to form a steep cusp, thus giving rise to characteristic differences in the two stellar populations that may be mirrored in the MW GC.  Because studies of the light excess and deficit in elliptical galaxies focus on major mergers, the scenario seems less relevant for a disc-dominated system like the Milky Way, which may have never experienced a major merger~\\citep{gilmore02}. A minor merger of the SMBH with an IMBH however cannot be ruled out. The presence of an IMBH in the Galactic centre has been previously considered as a possible vehicle for delivery of young stars into the GC~\\citep{hm03}, a mechanism for creation of hypervelocity stars~\\citep{baum06}, and for the growth of the SMBH~\\citep{pz06}. Indeed, the possibility that an IMBH with mass $\\lesssim 10^4 \\Msun$ is still lurking in the inner parsec of the GC cannot currently be totally excluded based on observations~\\citep{hm03,genzel10,rb04,gm09,ggm10}.  In light of the new observational evidence, which supports the notion that few to 10~Myr ago was a special period in the life of ${\\rm Sgr~A^*}$, as indicated by the relatively recent episode of star formation and increased energy output, we revisit the possibility that a minor merger could have triggered this epoch of enhanced activity. We suggest that the cumulative observational evidence favors the minor merger hypothesis relative to the scenarios that propose a steady state evolution or passive relaxation of the GC region. We present a theoretical scenario for one such minor merger in \\S~2 and discuss the implications in \\S~3.  \\section[]{Milky Way -- Satellite Merger Scenario}\\label{S_methods}  Here we examine the viability of the following scenario: at high redshift, a primordial satellite galaxy with a central IMBH begins to merge with a young Milky Way. As the satellite sinks toward the GC under the influence of dynamical friction it is tidally stripped and its orbit gradually decays toward the Milky Way disc plane~\\citep{qg86,callegari11}. The satellite perturbs previously stable gas clouds in the inner Milky Way disc, driving gas inflow~\\citep{noguchi88,bh91,bh96,hq10} and compressing the gas to densities exceeding those necessary for massive star formation near the GC~\\citep{mh96,hq10}. The satellite galaxy is expected to be largely disrupted by the time it reaches the GC, leaving the IMBH spiraling in a dense gaseous and stellar environment.  In the context of this scenario we hypothesize that the IMBH reached the central parsec on the order of $\\sim$10 Myr ago. A fraction of perturbed gas that did not form stars accretes onto the Milky Way's SMBH~\\citep{hq10}, injecting massive amounts of energy into the surrounding medium and giving rise to the Fermi bubbles~\\citep{su10,zubovas11}.  Once gravitationally bound, the IMBH-SMBH binary orbit tightens via three-body interactions with surrounding stellar background, scouring the old stellar population to form a central core~\\citep{merritt10}.  Finally, the binary coalesces after emitting copious gravitational radiation~\\citep{pm63}.  In the context of this hypothetical scenario we use the new GC observations to constrain the initial masses of the satellite and Milky Way galaxies ($M_{\\rm sat}$ and $M_{\\rm MW}$), the mass deficit in the late-type stellar population ($M_{\\rm def}$), the IMBH mass ($M_{\\rm IMBH}$), as well as the amount of gas inflow into the GC triggered by the inspiral of the satellite galaxy. The properties of the satellite bound to reach the inner disc of the Milky Way and deliver its IMBH to the GC must satisfy several criteria: 1. it should be light enough not to disrupt the Galactic disc, 2. it should be sufficiently massive in order for the dynamical friction to operate efficiently and deliver it to the GC within a Hubble time, and 3. its potential well should be sufficiently deep to sustain tidal stripping by the Milky Way. We therefore focus on constraining the most plausible scenario given the current understanding of the processes involved.  Our approach is, out of necessity, semi-analytical in nature. While advanced cosmological nbody simulations are capable of modeling the accretion of a low-mass satellite galaxy onto cosmologically growing Milky Way halo, there are a number of physical processes important to our model that these simulations cannot capture. For example, we will capture the effect of the Milky Way disc, bulge and SMBH, and will account for the stabilizing effect of the IMBH within the satellite. In our model, we also include the critical effects of 3-body scattering and gravitational wave emission, both of which are beyond the reach of a cosmological nbody simulation.     \\subsection{Properties of the Progenitor Milky Way}\\label{SS_gal} Beginning the merger at high redshift is advantageous in three respects.  First, at this early epoch, it is reasonable to assume that the proto-Milky Way was surrounded by primordial satellite galaxies capable of housing a central seed black hole~\\cite[e.g.,][]{rg05,gk06,wa08,micic11}.  Second, at this stage in its growth, the Milky Way would have been smaller, less massive, and more gas-rich than it is today, thus decreasing the time required for the satellite to sink to the galactic centre via dynamical friction. Finally, the orbits of infalling satellites are more radial at high redshift, which further shortens the merger time-scale~\\citep{wetzel11}. It should be noted that, while the remainder of this work posits that the satellite is accreted at redshift 8, this is by no means a unique solution.  To determine the properties of the Milky Way at this epoch, we assume that it grows according to the exponential halo model from \\citet{mcbride09}: \\begin{equation} M(z) = M_{z=0}(1+z)^{\\beta}{\\rm exp}\\left(-\\ln2\\frac{z}{z_{f}}\\right), \\label{eqn:halogrowth} \\end{equation} \\noindent where $M_{z=0}$ is the current halo mass and $z_f$ is the formation redshift, defined as the redshift at which the halo has grown to half its current mass. Adopting the properties for the Milky Way at $z_{f}=1$ as $M_{z=0}=2\\times10^{12} \\Msun$ and $\\beta=0.25$, the Milky Way's mass at $z=\\newredshift$ can be estimated to be $M_{\\rm MW} = \\newmwminit \\Msun$. Studies of cosmological N-body simulations have found that at the redshift considered, the concentration of dark matter haloes is very weakly dependent on mass \\citep{zhao03,gao2008,klypin11}. Following the methods outlined by \\citet{prada11}, we find that at $z=8$ a halo of this mass will have a concentration of $c(z=8)\\sim6$. The halo virial radius in a LCDM cosmology is defined as the radius where the mean enclosed density is 96 times the critical density of the universe, $\\rho_{\\rm crit}$. With the definition of  $\\rho_{\\rm crit}$:  \\begin{equation} {\\rho_{\\rm crit}} = {{{3 H_0^2}\\over {8 \\pi G }} \\left[\\Omega_{\\Lambda}+ (1+z)^3 \\, \\Omega_{m} \\right]},\\label{eqn:rhocrit} \\end\t{equation}  \\noindent where $\\Omega_{\\Lambda}=0.73$ is the fraction of energy density in the universe in vacuum energy, while $\\Omega_{m}=0.27$ is the fraction of energy density in the universe in matter, and z is the redshift. We find that the progenitor MW halo has a virial radius of $\\sim 6$ kpc. This implies the progenitor Milky Way halo will have a density at 10 pc of $\\newmwrho \\Msun/$pc$^{3} \\sim 10^6\\rho_{\\rm crit}$.  It is important to note that while Eqn. \\ref{eqn:halogrowth} assumes a single, smoothly growing Milky Way halo, at these high-redshifts, mergers with other massive haloes are very common, and the halo grows in a step-wise fashion.~\\citep{diemand07}. Indeed, the entire picture of a single, virialized progenitor Milky Way halo is not strictly correct, and the `Milky Way' at this redshift is more likely a set of several haloes, many of which have not yet decoupled from the Hubble flow to allow turnaround and collapse into a single virialized structure. Consequently, our assumption of a virialized NFW halo at the accretion redshift ($z=\\newredshift$) must be recognized as an approximation made due to the limits of a semi-analytic approach.  \\subsection{Finding the Culprit Satellite}\\label{SS_orb}  Broadly, we identify possible culprit satellites by integrating the orbits of infalling haloes within an analytic, but evolving Milky Way potential. As both the satellite and Milky Way evolve, we search for the satellites that reach the Inner Lindblad Resonance (ILR) at 150 pc roughly 10 Myr ago after losing over 95\\% of its initial orbital angular momentum. Of the satellites that survive until they plunge through the  ILR, we preferentially select those that retain enough mass to perturb the gas there. The culprit satellite is characterized by the mass, radius and concentration, as well as the energy, angular momentum, infall radius and merger redshift of the orbit. We elaborate on the procedure below.  We adopt a merger redshift of  $\\sim\\newredshift$. In order to deliver the IMBH to the GC a mere 2$-$7 Myr ago, the proposed merger redshift implies that the satellite orbit decayed over a time-scale of about \\newdecaytime{ }Gyr. At such a high redshift, the IMBH and satellite had very little time to evolve before being accreted by the Milky Way, making the pair a ``fossil'' of the dark ages before reionization \\citep{rg05,gk06}.  We rely on cosmological N-body simulations to constrain the initial conditions of the orbit. These inform us that at the present epoch, satellites are preferentially accreted on very eccentric orbits, with a distribution peak at about $e=0.85$ \\citep{benson05,wang05,zentner05,khochfar06,ghigna98,tormen97}. At higher redshifts the satellite orbits are characterized by even higher eccentricities, albeit, in both cases the distribution peaks are broad. Seemingly independent of redshift, a typical satellite is accreted at the virial radius with a total velocity, $|\\vec{v}_{\\rm sat}|=1.15 v_{\\rm vir}$ ($v_{\\rm vir}$ is the circular velocity at the virial radius of the primary galaxy) that marks it as barely bound \\citep{benson05,wetzel11}. Motivated by these results, we select an orbit that has $|\\vec{v}_{\\rm sat}|=1.15 v_{\\rm vir}$ at the virial radius of the primary and an eccentricity of $\\newecc$, consistent with expectations for the eccentricity distribution peak at $z=\\newredshift$ \\citep{wetzel11}.   Starting with the above total velocity and eccentricity, we calculate the orbital decay for a range of satellite masses placed at the virial radius of the primary. For a given initial position at the virial radius, the azimuthal and radial components of the satellite's initial velocity within the orbital plane are calculated in terms of the eccentricity ($e$) and total velocity ($|\\vec{v}_{\\rm sat}|$) as: \\begin{equation} v_{\\phi} = \\frac{v_{vir}}{|\\vec{v}_{\\rm sat}|}\\sqrt{\\frac{GM_{\\rm MW}}{r_{\\rm vir}}(1-e^2)} \\mbox{ and } \\\\ v_{r} = \\sqrt{|\\vec{v}_{\\rm sat}|^{2}-v_{\\phi}^{2}}. \\end{equation}  We adopt an analytic model of the Milky Way that includes a central SMBH, Miyamoto-Nagai thin disc~\\citep{mn75}, a spherical Hernquist bulge~\\citep{hernquist90}, and an NFW halo~\\citep{nfw97}. To mimic a young Milky Way, we  use  Eqn. \\ref{eqn:halogrowth} to set the halo mass. We set the virial radius using the mass and the critical density at the starting redshift in Eqn. \\ref{eqn:rhocrit}, and we initialize the concentration using \\citet{prada11}. We assume that the mass and size of the baryonic components change in the same way as the halo does; this is not true in detail, but allows us to convert the known present-day Milky Way parameters to the starting redshift.  Our current Milky Way mass model is similar to analytic models best-fit to rotation curve data~\\citep[e.g.][]{wid05, den98} $z=0$ disc mass is $5\\times10^{10} \\Msun$, the disc scale length is $3$ kpc and, and the disc scale height is $300$ pc.For the bulge, we set a current epoch bulge mass of $8\\times10^{9} \\Msun$ and scale length of $0.7$ kpc.    We integrate the orbits using a fourth-order Runge-Kutta method to step the satellite's position and velocity forward in time. At each timestep, we adjust the analytic Milky Way model using the method described above. We calculate the acceleration of the satellite due to this evolving analytic potential, and we  include  Chandrasekhar dynamical friction~\\citep{chandra43}, as well as mass loss from the satellite due to tidal stripping and disc shocks. The acceleration due to dynamical friction is calculated in the uniform density limit as \\begin{eqnarray} \\left(\\frac{d\\vec{v}_{\\rm sat}}{dt}\\right)_{\\rm fric} &=& -\\frac{4\\pi\\ln \\Lambda G^{2}M_{\\rm sat}\\rho_{\\rm MW}}{|\\vec{v}_{\\rm sat}|^{3}}\\times \\cdots\\nonumber\\\\ &&\\times\\left[{\\rm erf}(\\chi)-\\frac{2\\chi}{\\sqrt{\\pi}}e^{-\\chi^2}\\right] \\vec{v}_{\\rm sat}, \\label{eqn:accfric} \\end{eqnarray} where $M_{\\rm sat}$ is the mass of the satellite, $\\rho_{\\rm MW}$ is the density of the Milky Way at the satellite's position, $\\ln\\Lambda = \\ln\\left[1+ (M_{\\rm MW}/M_{\\rm sat})^2\\right]$ is the Coulomb logarithm, $\\chi=|\\vec{v}_{\\rm sat}|/\\sqrt{2}\\sigma$, and $\\sigma=\\sqrt{GM_{\\rm MW}/2R_{\\rm MW}}$ is the average velocity dispersion of the Milky Way halo.   At each step in the orbit, we calculate the local density of the Milky Way and we tidally strip the satellite to the Roche radius, where the density of the satellite is equal to the Milky Way background. We also model mass loss from disc shocking by removing \\begin{equation} \\Delta M_{\\rm shock} = \\frac{5}{3}\\frac{4}{GM_{\\rm sat}v_{\\rm sat,z}^{2}}\\left(\\frac{dv_{\\rm sat,z}}{dt}\\right)_{\\rm disc}^{2} \\label{eqn:shock} \\end\t{equation} from the satellite's mass each time it passes through the Milky Way disc \\citep{gnedin97}. We neglect the stellar component of the satellite, since the baryon content of such low mass satellites is relatively uncertain, but likely to be very small \\citep{gnedin00,sg07,ricotti08}.   We find that the most likely culprit is a satellite with a mass of $M_{\\rm sat} \\approx \\newsatminit \\Msun$. Modeling the satellite dark matter profile as an NFW halo, its corresponding concentration parameter at this redshift is about \\newsatconcen~\\citep{prada11}, making the satellite's central density within the inner 10 pc $\\sim\\newsatrho \\Msun/$pc$^{3} \\sim 4 \\times 10^5 \\,\\rho_{crit}$ or $\\sim 2\\times10^4$ times the Milky Way's density at the virial radius. Including an IMBH in our satellite model would deepen its central potential and could aid in delivering the satellite core to the centre of the MW intact, although we did not include this effect in our calculations.   By the time the satellite has reached the inner 100 pc, it will have lost most of its mass, with $\\sim \\newsatmfinal \\Msun$ remaining.  Without direct hydrodynamic simulations, it is difficult to say how much damage this IMBH-embedded satellite core could do to the gas-rich inner Milky Way.  In general, we expect the satellite to perturb the gas in the galactic centre, torquing it and transporting angular momentum through narrow resonances~\\citep{gold79}; the classical rate of gas inflow from this process is proportional to the strength of the perturbation squared. However, when the system has a significant asymmetric perturbation, the orbits begin to cross one another and gas piles up in shocks~\\citep{papa77}. In this case, the radial inflow rate of gas from a global perturbation is linearly proportional to the strength of perturbation~\\citep{hq11}, and numerical simulations find the shocks induced by even a few $\\%$ perturbation can destabilize the gas and drive gas inflow \\citep{bh91,bh96,mh96,hq10}. To estimate the perturbation a $\\sim \\newsatmfinal \\Msun$ satellite core could exert on the gas accumulated in a ring at the Inner Lindblad Resonance (ILR) of the Milky Way, we refer to ~\\citet{vw00}, which explores the perturbation strength induced by galaxy flyby encounters.   Using linear perturbation theory, ~\\citet{vw00} find that a flyby with a mass ratio of 10 and a pericentre at the half-mass radius will induce a strong perturbation in the density of the primary galaxy of order unity.  Since the mass ratio of the inner Milky Way ($\\sim10^{8} \\Msun$) to the satellite remnant is 1000 \\citep{lindqvist92}, we expect a perturbation of the order $|a|\\sim0.01$ in the surface density. The linear relationship between gas inflow and perturbation amplitude derived by \\citet{hq11}, \\begin{equation} \\frac{dM_{\\rm gas}}{dt}=|a|\\Sigma_{\\rm gas}R^{2}\\Omega, \\end{equation} can then be used to gauge the expected amount of gas inflow. Setting the perturbation amplitude to $|a|\\sim0.01$, the radius to $R=150$ pc (ILR), the rotation frequency to $\\Omega(R)=v_{circ}(R)/R=0.62$ Myr$^{-1}$ \\citep{stark04}, and the gas surface density to $\\Sigma_{\\rm gas}=500\\Msun/$pc$^{2}$ based on observations of other barred galaxies \\citep{jogee05} and of the molecular ring in the MW GC \\citep{molinari11}, yields a gas inflow rate of $\\sim7\\times10^4 \\Msun/$Myr. Assuming this inflow rate over $\\sim10$ Myr, we find that this satellite should be able to drive a net inflow of $\\sim10^{6}\\Msun$ of gas from the ILR.   \\subsection{Late-Type Stellar Mass Deficit and IMBH Mass}\\label{SS_mdef} If the core in the distribution of late-type stars at the GC \\citep{buchholz09,do09} was scoured out by an IMBH-SMBH binary~\\citep{preto11,gm12}, the amount of stellar mass missing from the GC can be used to constrain the mass ratio of the black hole binary~\\citep{mm01,gm08,merritt06}. To determine this mass deficit we compare the stellar distribution inferred from observations with that expected for a dynamically relaxed system without a core. In terms of the number density of the late-type stellar population, the core can be represented by a broken power law \\begin{equation} \\label{dens_fin} n_{f}(r) = n_{0}\\left(\\frac{r}{r_{0}}\\right)^{-\\gamma_{i}} \\left[1+\\left(\\frac{r}{r_{0}}\\right)^{\\alpha}\\right]^{(\\gamma_{i}-\\gamma)/\\alpha}, \\end{equation} \\noindent with $n_{0}=0.21$~pc$^{-3}$, $r_{0}=0.21$~pc, $\\gamma=1.8$, $\\gamma_{i}=-1.0$, and $\\alpha=4$ \\citep{merritt10}. We adopt this description in our analysis but note that in presence of strong mass segregation the slopes can be steeper \\citep{AlexanderHopman09,PretoAmaroSeoane10,Amaro-SeoanePreto11}. The observed distribution of stars outside of the $0.21$~pc core radius is consistent with the Bahcall-Wolf profile \\citep[$\\propto r^{-1.75}$,][]{bahcall76} of a relaxed system as it would have existed prior to scattering by the IMBH-SMBH binary.  We model the initial stellar cusp by extending the $r^{-1.8}$ profile to smaller radii: \\begin{equation} \\label{dens_in} n_{i}(r) = n_{0}\\left(\\frac{r}{r_{0}}\\right)^{-\\gamma}. \\end{equation} Assuming that the mass density profiles {\\it before} and {\\it after} the creation of the core are proportional to equations (\\ref{dens_in}) and (\\ref{dens_fin}) respectively, we calculate the mass deficit as the integrated difference between the initial and final (observed) profiles. We normalize the profile given by Eqn. \\ref{dens_fin} such that integrating it over the inner parsec yields $1.0\\pm0.5\\times10^{6} \\Msun$, the mass determined by \\citet{schodel09}, and obtain $M_{\\rm def} \\approx 2\\times10^{5} \\Msun$.  It should be understood that this mass deficit can only be treated as an estimate. Although this calculation assumes the best fit core radius ($r_{0}$) and inner slope ($\\gamma_{i}$) from \\citet{merritt10}, the fit was not excellent ($\\tilde{\\chi}^2>17$) and the estimated mass deficit is highly dependent on these parameters. In addition, this calculation assumes that core size (and therefore the mass deficit) has not changed significantly over time. This is consistent with a core scoured recently enough ($\\sim10$ Myr) that relaxation has not yet had enough time to fill in the core \\citep[$\\sim10$ Gyr;][]{merritt10}. However, it is also possible that the core is an evolved system; this implies a larger core, more massive IMBH, and more dramatic scouring event in the distant past. In this case, the creation of the core would have been unrelated to the creation of the young GC stars or Fermi bubbles.  N-body merger simulations studying the relationship between the ratio of total stellar mass ejected to binary mass, $M_{\\rm def}/(M_1+M_2)$, and binary mass ratio, $q = M_1/M_2$, have not yet been carried out for the mass deficit calculated here. In order to relate the two we use a semi-analytic formalism describing the interaction of massive black hole binaries with their stellar environment \\citep{sesana08} to place the upper and lower limits on the mass of the IMBH based on $M_{\\rm def}$ inferred from observations.  It has been shown by numerical simulations \\citep{baum06,matsu07} and semianalytic models \\citep{sesana08}, that an IMBH inspiralling in a stellar cusp surrounding a central SMBH starts to efficiently eject stars at a separation $a_0$, where the stellar mass enclosed in the IMBH orbit is of the order of $2M_2$. The ejection of bound stars causes an IMBH orbital decay of a factor of $\\approx 10$, excavating a core of radius $r_0\\approx 2a_0$ in the central stellar cusp, resulting in a mass deficit about $3M_2$ \\citep[see][for details]{sesana08}. Such orbital decay is in general insufficient to bring the IMBH in the efficient gravitational wave (GW) emission regime, unless its eccentricity grows to $>0.9$ during the shrinking process. It is also the case in this picture that the mass of the inspiralling IMBH inferred for a given mass deficit strongly depends on the eccentricity evolution of its orbit. In what follows, we consider both the high and low orbital eccentricity scenario and use them to place a bound on the plausible range of IMBH masses.  If the eccentricity grows efficiently, the IMBH depletes the central cusp, forms a core of a size $\\approx 2a_0$, and merges due to GW emission on a time scale of only $1-10$~Myr~\\citep{sesana08}. For a stellar distribution described by an isothermal sphere outside of the radius of influence of the SMBH, $a_0 = 2q^{4/5}$pc. Adopting the core radius of $r_0=2a_0=0.21$ pc, we find $q=0.02$, and an upper limit on the mass of the IMBH, $M_2 = 8\\times10^4 \\Msun$. In this case, the mass evacuated from the stellar cusp by the IMBH is of the order of $3M_2$~\\citep{sesana08}, i.e., $\\approx2.5\\times 10^5$M$_\\odot$, consistent with the stellar mass deficit measurement in the GC.  Alternatively, if the IMBH eccentricity does not grow significantly during the bound cusp erosion, further scattering of stars replenishing the binary loss cone is needed in order to evolve from separation of $a_0$ to the GW regime. Therefore, a circular orbit regime can be used to establish a lower limit on the mass of the IMBH, for a given mass deficit indicated by observations. We assume that in this case both $r_0$ and $M_{\\rm def}$ created in the cusp erosion phase are small (we justify this assumption below). In this scenario, the final $r_0$ and $M_{\\rm def}$ are reached as a consequence of the diffusion of the stars from the edge of the small core into the loss cone of the binary. The ejections of each star carry away an energy of the order $(3/2)G\\mu/a$~\\citep{quinlan96}, where $\\mu=M_1\\,M_2/M$. We compute $M_{\\rm def}$  by imposing: \\begin{equation} \\frac{3}{2}\\frac{G\\mu}{a}dM_{\\rm def}=\\frac{GM_1M_2}{2}d\\frac{1}{a} \\end{equation} to get \\begin{equation} M_{\\rm def}=\\frac{M_1+M_2}{3}{\\rm ln}\\frac{a_i}{a_f}, \\end{equation} where $a_{\\rm i}$ is the hardening radius of the binary (radius at which the scattering of unbound stars becomes effective) and $a_{\\rm f}$ is the separation at which the GW emission becomes efficient. Using equations (19) and (20) in~\\cite{sesana10} to express $a_{\\rm i}$ and $a_{\\rm f}$, it follows that, \\begin{equation} M_{\\rm def}=\\frac{M_1+M_2}{3}\\,{\\rm ln}\\frac{500\\,q^{4/5}}{F(e)^{1/5}}, \\end{equation} where $F(e)=(1-e^2)^{-7/2}(1+73/24\\,e^2+37/96\\,e^4)$. Assuming for the purpose of this estimate that the binary remains circular throughout its evolution and imposing $M_{\\rm def}=2\\times 10^5 \\Msun$, we find $q=5\\times10^{-4}$ and a lower limit on the mass of the IMBH is $M_2=2\\times10^3 \\Msun$ \\footnote{Both numerical simulations and semi-analytic models however suggest that the eccentricity in the cusp erosion phase grows to $>0.9$, in which case $F(e)>1000$ and $q>5\\times 10^{-3}$, i.e., $M_2>2\\times10^4 \\Msun$.}.  An IMBH of such mass, would excavate a core of $\\approx 0.01$ pc, causing a mass deficit of $\\sim 3M_2=6\\times10^3 \\Msun$ in the bound scattering phase and thus, justifying our earlier assumption that the diffusion of stars into the loss cone is the primary process that shapes the properties of the core in this case.  Note that in the circular orbit scenario the time scale for the inspiral of the IMBH towards the GW regime is determined by the unknown rate of diffusion of the stars into the loss cone of the binary. Hence, depending on the time scale of relaxation processes this process could in principle lead to the IMBH-SMBH binary ``hangup'', i.e., a long lived ($> 1$~Gyr) binary configuration at separation $<r_0$ -- tantamount to the classical ``final parsec'' problem~\\citep{BegelmanEtAl80}. It is however possible that the binary will not stall in our specific case. The galactic centre in this phase will be described by a strongly perturbed, non-axisymmetric potential which allows stars to scatter into the loss cone efficiently \\citep{merritt04,berzcik06,perets08,khan11}.  Moreover, the orbit will occur in a relatively gas-rich environment, which can further aid the decay of the binary \\citep{escala05,dotti07,cuadra09}. Finally, any extra stars brought in by the satellite would help the binary decay~\\citep[see][]{Miller02}. Even under the assumption of a circular orbit, an efficient coalescence can occur on a time-scale of 10~Myr.  This analysis suggests that the observed mass deficit and core size are consistent with the IMBH mass in the range $2\\times 10^3 \\Msun <M_2<8\\times 10^4 \\Msun$, whereas the efficient eccentricity growth found in N-body simulations and semi-analytic models favor $M_2\\gtrsim 10^4 \\Msun$. Within this range, the time scale for the IMBH to create a core and merge with the SMBH can be as short as few Myr. On the other hand, a possibility that an IMBH may be still be lurking in the GC is not completely ruled out. We discuss the consequences of the latter scenario in the context of the observational constraints on the presence of a second black hole in the Galactic centre in \\S~\\ref{S_discuss}.  It is useful to consider whether a satellite galaxy with an initial mass of $M_{\\rm sat}\\sim \\newsatminit \\Msun$ can host a $\\gtrsim 10^4 \\Msun$ IMBH.  While there are no observational constraints for galaxies or black holes of this mass range, there are three leading theories for IMBH formation at high redshift: `direct collapse' of metal-free, low angular momentum gas into a $10^3-10^6 \\Msun$ black hole~\\citep{loeb94,begelman08}, an unstable supermassive star that collapses into a $10^2-10^5 \\Msun $ black hole \\citep{colgate67,quinlan87, baumgarte99}, or a Population III star, which would leave behind seed black holes of $\\sim1-10^3 \\Msun$ between redshift 30$-$12~\\citep{madau01,bromm02,wa08,clark11}. Even if the IMBH in our satellite started as a low mass Pop III seed in a somewhat turbulent environment with a mass of $\\sim5 \\Msun$~\\citep{clark11}, it is plausible that it would reach the IMBH mass proposed here through a combination of gas accretion and black hole mergers~\\citep{KHB10}. In such a satellite galaxy, it would require less than one percent of the gas to accrete onto a low mass seed to form the IMBH $\\gtrsim 10^4 \\Msun$.  Note that the massive seeds produced in a direct collapse typically favor more massive haloes than the one we have proposed as our culprit. This is because metal-free gas collapses most efficiently in haloes with $T_{\\rm vir} > 10^4$ K, corresponding to $M_{\\rm vir} > 10^8 \\Msun [(1+z)/10]^{3/2}$ \\citep{bromm03}. In the context of the merger hypothesis choosing a slightly more massive satellite would push the accretion redshift closer to the present day, and as long as the resulting satellite merger is still a minor one, this does not significantly affect the outcome of our scenario.   \\subsection{Inflow of Gas and Gamma-ray Bubbles}\\label{SS_bubble} As noted in \\S~\\ref{SS_gal} the inspiral of a satellite galaxy can cause the inflow of a significant amount of gas towards the centre of the Galaxy~\\citep{noguchi88,bh96,mh96,cox08}.  One fraction of this gas could have given rise to the star formation in the Central, Arches, and Quintuplet clusters, which marked the epoch between 2$-$7~Myr ago in the central 50~pc of the Milky Way. All three clusters contain some of the most massive stars in the Galaxy and have inferred masses of $\\sim10^4 \\Msun$ \\citep{figer08}. Assuming a ``standard'' star formation efficiency of 10\\%~\\citep{ry99}, it follows that the amount of gas necessary to produce the stellar population of the three clusters is a ${\\rm few} \\times 10^5 \\Msun$. Note that a sequence of strongly compressional events during the satellite-Milky Way merger could have given rise to a higher efficiency of star formation~\\citep{dimatteo07}, in which case the estimated mass of the gas represents an upper limit.  In this merger scenario, the remainder of the perturbed gas that did not form stars would be channeled towards the central parsec \\citep{loose82}, and the fraction that is accreted into the SMBH could drive the energetic outburst of several Myr ago. The far-IR and millimeter observations indicate that $\\sim 10^4 \\Msun$ of the molecular gas continues to reside in the circumnuclear disc within the central $\\sim1.5$~pc of the Galaxy \\citep[see][for review and references therein]{genzel10}. The maximum amount of the remnant molecular gas that has not been accreted onto the SMBH can also be estimated based on its expected gravitational effect on the orbits of the stars residing within the inner 0.5 pc. In this case, the requirement for stability of the stellar disc over its lifetime of 6~Myr poses a constraint on the mass of the molecular torus of $<10^6 \\Msun$ \\citep{subr09}.  On the other hand, the recent discovery of the two large gamma-ray bubbles extending from the GC above and below the galactic plane are compelling evidence of a relatively recent period of intense activity in the now quiet GC. The gamma-ray bubbles exhibit several striking properties: they are perpendicular and symmetric with respect to the plane of the Galaxy, have nearly uniform gamma-ray brightness across the bubbles, and well defined sharp edges \\citep{dobler10,su10}. The gamma-ray emission from the bubbles is characterized by the hard energy spectrum and is most likely to originate from the inverse Compton scattering of the interstellar radiation field on the cosmic ray electrons -- the same population of electrons deemed responsible for the diffuse synchrotron microwave radiation detected by the WMAP \\citep{finke04,df08}. The sharp edges of the Fermi bubbles are also traced by the X-ray arcs discovered in the ROSAT maps \\citep{snowden97}, suggested to be the remnants of shock fronts created by the expanding bubbles \\citep{su10,guo11}.  The morphology, energetics, and emission properties of the Fermi bubbles favor the explanation that bubbles were created in a strong episode of energy injection in the GC in the last $\\sim 10$~Myr that followed an accretion event onto the SMBH \\citep{su10}. Simulations by \\citet{guo11} indicate that the bubbles could have been formed by a pair of bipolar jets that released a total energy of $1-8\\times 10^{57}$~erg over the course of $\\sim$0.1$-$0.5~Myr between 1 and 2~Myr ago. This explanation for the Fermi bubbles implies that $\\sim 10^4 \\Msun$ of material must have been accreted onto the SMBH at nearly the Eddington rate, assuming the accretion efficiency of $10\\%$ \\citep{ss73,dl11}. Based on the range of models explored by \\citet{guo11} it is possible to estimate that the amount of mass processed in such jets (i.e., the mass of the gas that fills the jet cavities) is as small as $30 \\Msun$ and as large as $3\\times10^5\\Msun$.  This estimate, together with the gas that formed stars, the gas accreted onto the SMBH and the gas processed by the jets allows us to put a constraint on the total gas inflow into the central $\\sim50$~pc of the Galaxy of $\\lesssim 10^6 \\Msun$, consistent with the amount expected from the perturbation analysis of the stability of the ILR gas in the Milky Way.  \\section[]{Discussion}\\label{S_discuss}  \\subsection{How rare are satellite merger events?} We propose that the timeline began about \\newdecaytime{} Gyr ago, when the proto Milky Way accreted a small satellite dark matter halo at the time when their haloes were physically closer and less massive. The satellite orbit decayed slowly and only reached the GC a few million years ago, after having been stripped of most its mass. The thinness of the Milky Way disc has often been used as an argument against a recent minor merger~\\citep{qhf93,snt98,vw99}; however, the proposed satellite is so minor, particularly by the time the orbit decays to 10 kpc, that the thin disc could have survived unscathed \\citep{toth92,walker96,tb01,hopkins08,hopkins09}.   Using the Extended Press-Schechter formalism (EPS)~\\citep{bond91,bower91,lc93,pch08}, we can estimate the number of satellite accretion events a typical Milky Way mass galaxy will undergo. We determined this on the basis of 100 realizations of an EPS merger tree that resulted in a base halo of $2\\times 10^{12} \\Msun$ at z=0. In this calculation, we assumed WMAP5 parameters and summed over the haloes in the $M_{\\rm sat}=10^7-10^9 \\Msun$ mass range that merged with the main halo from z=7 to z=0.We found a mean of 1745 such satellite accretion events, with a standard deviation of 425. However, about half of these accretions occur after $z=1$ --- and are unlikely to have made it to the GC by $z=0$. We confirmed that this number of satellite accretion events is consistent with expectations from the cosmological simulations by comparing to one of our N-body simulations of a 50 Mpc$^3$ volume described in ~\\citet{shb11}.       \\begin{figure*} \\includegraphics[height=7in,angle=270]{fig1.eps} \\vspace{-0.6in} \\caption{Distribution of accreted low mass satellites at the present day. The inner 40 kpc region of the Milky Way disc is shown in greyscale with current accreted satellite positions overplotted. The color maps to the surface brightness, and the relative size corresponds to the tidal radius of the satellite. Note that the circle size is not to scale and that none of these satellites would be observable above the background. Also note that satellites that have merged with SMBH or completely disrupted are not plotted.} \\label{fig:conveyer} \\end{figure*} Figure~\\ref{fig:conveyer} illustrates one realization of the current distribution of accreted satellites in the mass range $M_{\\rm sat}=10^7-10^9 \\Msun$. We used the EPS technique described above to define the number of accreted satellites in this mass range at each integer step in redshift from z=7 to z=0. To define the orbit of each satellite as it is accreted, we randomly selected from the energy and angular momentum distributions at each redshift using the expressions $7-9$ in~\\citep{wetzel11}.  As in section 2, we integrated orbits of the satellites from the accretion epoch to the present day, scaling the Milky Way mass and size to the redshift of accretion using Eqn.~\\ref{eqn:halogrowth}. Figure~\\ref{fig:conveyer} shows the inner 40~kpc of the current-day Milky Way; approximately 85\\% of the accreted satellites are at separations larger than 40~kpc, and only 5 reached the GC and merged with the SMBH. We estimate the surface brightness of the satellites assuming that the baryons are confined to a radius within the dark matter halo ten times smaller than the satellite virial radius. We infer the initial star fraction from \\citet{rg05} and assume a total mass-to-light ratio of $\\sim300$ \\citep{strigari08} for the bound stars that remain after tidal stripping.    While it is very clear that not all of the small satellites can reach the GC, what fraction does is a question of some subtlety. Galaxy merger time-scales cited in the literature, particularly for the small mass ratios considered here, span a wide range. The key to the uncertainties is the treatment of dynamical friction: most semi-analytic works, including this one, rely on the dynamical friction formalism as described by Chandrasekhar (1943) but change the Coulomb logarithm to account for inhomogeneous or anisotropic systems~\\citep{pen04,just05}, or to include mass loss~\\citep{tb01,vw99}. This approach has been shown to underestimate the decay time in pure dark matter simulations \\citep{colpi99,bk08}. On the other hand, the presence of gas can dramatically decrease the orbital decay time of a satellite by efficiently dissipating its orbital energy throughout the system \\citep{ostriker99,sanchez99} -- thus, making the Chandrasekhar formula a significant overestimate. Compounding the issue, linear perturbation theory and limited N-body experiments indicate that resonant heating caused by orbits in the satellite galaxy that are commensurate with the orbit of the satellite about the GC can enhance mass loss and can change the angular momentum of the orbit in non-trivial ways~\\citep{weinberg97,choi09}. Although the Milky Way has likely accreted over a thousand of these small satellites, it is uncertain how often they reached the galactic centre. It is however plausible that the GC has experienced a handful of these accretion events spread over its lifespan. Despite the uncertainty that arises in the mass of our culprit satellite due to the imprecise dynamical friction time-scale, the scenario itself remains viable, because other constraints on satellite mass (satellite evaporation, disc disruption) are flexible so long as the merger time-scale remains less than the age of the universe.    \\subsection{Hypervelocity stars and stellar core}  While the properties of the newly formed stars and perturbed gas were dictated by the accreted satellite, the IMBH was responsible for carving out the old stellar population. As a gravitationally bound binary IMBH-SMBH formed and decayed, it scoured out $2 \\times 10^5 \\Msun$ of the relaxed old and initially cuspy stellar population. Many of these stars could have been ejected from the GC as hypervelocity stars \\citep[HVSs;][]{brown05,baum06}, though most may simply have received enough energy to traverse the inner parsec. Simulations of IMBH-SMBH binaries in stellar environments indicate that HVSs are created in a short burst which lasts only a few Myr in case of a $\\sim 10^4 \\Msun$ IMBH \\citep{baum06, sesana08}.  In the context of our picture we predict that this event created $\\sim 10^3$ hypervelocity stars that, if they were ejected at about 1000~${\\rm km\\,s^{-1}}$ \\citep{baum06}, ought to lie $\\sim10$ kpc from the GC today. It is worth noting that about a dozen of HVSs observed in the Galactic halo thus far have travel times that span 60$-$240~Myr and appear to be consistent with a continuous ejection model \\citep{brown08,brown09,tillich09,irrgang10} and not with the IMBH-SMBH binary picture \\citep{brown08,sesana08}. Along similar lines, the spatial and velocity distribution of the current observed HVSs seem to be inconsistent with a IMBH-SMBH slingshot origin \\citep{sesanaetal07}.  The large size of the observed GC core, $r_{0}=0.21$~pc, could be seen as a challenge to any scenario involving 3-body scattering, since state of the art high resolution direct N-body simulations that modeled the ejection of hypervelocity stars from a SMBH-IMBH binary in the galactic centre never generated a core larger than 0.02 parsecs \\citep{baum06}. However, there are several effects that could conspire to cause the simulated core size to be a lower limit. First, the mass of the simulated SMBH in \\citet{baum06} is $3\\times 10^6 \\Msun$, which would eject fewer stars than somewhat more massive Milky Way SMBH. Second, the density profile was sharply curtailed by a factor of $(1+r^5)$ in order to minimize the number of stars far from the SMBH; this makes the spatial distribution of stars in the simulated nuclear star cluster more centrally peaked relative to that in the GC, which can also result in a smaller core. In general, though, it is important to note that the size of the scoured core is a property that sensitively depends on the density, the eccentricity, and kinematic structure of the GC or on assumptions in the model used to represent it.  \\subsection{Has IMBH-SMBH binary merged?}  We now return to the question whether the IMBH-SMBH binary has already merged or whether the IMBH could still be lurking in the GC. As discussed in \\S~\\ref{SS_mdef}, the N-body and semi-analytic modeling of the GC favor the evolutionary scenarios in which the inspiral and coalescence of the SMBH with a $M_2\\gtrsim 10^4 \\Msun$ IMBH is relatively efficient. Moreover, there is currently no empirical evidence for a second black hole in the central parsec. In order to be consistent with the observations, the IMBH present in the GC would have to have a mass $\\sim 10^3 - 10^{4.5}\\Msun$ and be either very close ($\\leq 10^{-3}$ pc) or at $>0.1$ pc from the SMBH \\citep{rb04,gm09,ggm10,genzel10}. An IMBH in this mass range that reaches a separation of $10^{-4}$ pc would merge with the SMBH in less than 10~Myr due to the emission of GWs, thus severely restricting the amount of parameter space where the IMBH and SMBH can exist in a long lived binary configuration. Nevertheless, given the uncertainties in the binary mass ratio, eccentricity, and the structure of the initial stellar cusp, the presence of an IMBH in the GC cannot be entirely ruled out at this point.  If on the other hand, the IMBH and SMBH coalesced several million years ago, one possible signature of this event could be a SMBH recoil caused by the asymmetric emission of GWs \\citep{per62,bek73}. Current astrometric observations of the reflex motion of the SMBH put strong constraints on the allowed recoil velocity; the SMBH cannot have velocity with respect to the Central cluster larger than $3.5$ km/s (within $1\\sigma$ error), at the distance of the GC \\citep{yelda10}. Similarly,  \\citet{rb04} constrain the peculiar motion of Sgr~A* in the plane of the Galaxy to $-18\\pm7$ km/s and perpendicular to the Galactic plane to $-0.4\\pm0.9$ km/s, where quoted uncertainties are $1\\sigma$ errors. There is however a caveat with respect to the interpretation of the SMBH reflex motion: if the reference frame in which the reflex motion is measured is based on the nearby gas and stars bound to the SMBH, the resulting relative velocity of the SMBH will be zero because in this case, the stars and the gas move together with the SMBH as long as their orbital velocity is higher than the that of the reflex motion. The radio and near-infrared reference frames in \\citet{yelda10} are defined based on the nearby stars orbiting around the SMBH and are thus a subject to this caveat. The measurement of \\citet{rb04} is however carried out in the reference frame defined by the extragalactic radio sources and can be used to test the recoil hypothesis.  For $10^4 \\Msun$ IMBH the black hole merger can give rise to a modest recoil velocity of about $80$ m/s, assuming that the IMBH is not spinning rapidly. The recoil velocity magnitude in this case scales as $\\propto q^2$ \\citep{campanelli07,baker08}, thus implying that the coalescence of the SMBH with a slowly spinning IMBH more massive than $1.5 \\times 10^5 \\Msun$ can be ruled out based on larger of the observational constraints, as long as damping of the recoil motion of a remnant SMBH is inefficient on the time scale of several million years. More stringent constraints on the mass of the IMBH, based on the motion of the SMBH perpendicular to the Galactic plane, can be placed given the (unknown) orientation of the orbital plane of the binary before the merger in addition to the binary mass ratio and the spin vector of the IMBH.   \\subsection{Orientation of the SMBH spin axis}\\label{SS_spin}  The nearly perpendicular orientation of the spin axis of the SMBH to the Galactic disc plane, indicated by the orientation of the observed gamma-ray bubbles and jets in simulations of \\citet{guo11}, implies that the evolution of the SMBH spin has been determined by accretion from the Galactic gas disc rather than random accretion events with isotropic spatial distribution. Such events would include tidal disruptions of stars and giant molecular clouds triggered by the satellite inspiral and a merger with the satellite IMBH which orbital plane in principle may not be aligned with the plane of the Galaxy. It is thus interesting to consider whether a sequence of such accretion events can exhibit a cumulative torque on the SMBH sufficient to re-orient its spin axis, assuming that before the merger with a satellite galaxy it was perpendicular to the Galactic plane.  Consider first the effect of episodic gas accretion resulting from multiple tidal disruption events. \\citet{chen09,chen11} show that three-body interactions between bound stars in a stellar cusp and a massive binary with properties similar to the IMBH-SMBH considered here can produce a burst of tidal disruptions, which for a short period of time ($\\sim0.1$~Myr) can exceed the tidal disruption rate for a single massive black hole by two orders of magnitude, reaching $\\dot{N}\\sim 10^{-2}\\,{\\rm yr^{-1}}$. This implies that in the process of the IMBH inspiral the SMBH could have disrupted $\\sim 10^3$ stars. A key element in this consideration follows from the finding by \\citet{np98} and \\citet{na99} that the orientation of the spin axis of a SMBH is very sensitive to the angular momentum of the accreted gas: namely, accretion of a mere few \\% in mass of a SMBH can exert torques that change the direction but not the magnitude of the spin of a black hole. Because each in a sequence of random accretion events imposes an infinitesimal change in the orientation of the SMBH spin axis, collectively they can cause the spin axis to perform a random walk about its initial orientation. Thus, the magnitude of the effect scales with the number of disrupted stars and their mass as $\\sim \\sqrt{N}\\,m_*$. Since this is much less than few percent of $M_1$, the cumulative effect of tidal disruption events on the orientation of the spin axis of the SMBH will be negligible.  This conclusion is reinforced by an additional property of post-tidal disruption accretion discs: they are compact in size and usually confined to the region of a size ${\\rm few}\\times r_t$, where $r_t \\approx r_*\\,(M_1/m_*)^{1/3}$ is the tidal disruption radius of a star and $r_*$ is the stellar radius \\citep{rees88}. Such small accretion discs effectively act as very short lever arms for torques acting on the spin axis of the SMBH, thus further reducing the efficiency of this process \\citep{np98}.  Similar conclusions can be reached about the tidally disrupted molecular clouds and  gas flows that plunge towards the SMBH on nearly radial orbits as a consequence of perturbations excited by the satellite galaxy. In section \\S~\\ref{SS_bubble} we estimated that the amount of mass accreted by the SMBH is $\\sim 10^4 \\Msun$. A modest mass, combined with the small circularization radius of the gas accretion disc is insufficient to cause a significant change in the SMBH spin orientation. Even ``accretion'' of a spinning IMBH is not expected to noticeably influence the spin orientation of the remnant SMBH. The large mass ratio of the binary ensures that the final contribution of the IMBH's spin and orbital angular momentum to the final spin of the SMBH is small, as long as the pre-merger SMBH has a moderate initial spin, $>{\\rm few}\\times 0.1$, in terms of the dimensionless spin parameter \\citep{barausse09}. Hence, coalescence with the IMBH would not have had a significant effect on the SMBH spin axis orientation.  In summary, the torques from the accretion of tidally disrupted stars, gas, and the IMBH in the aftermath of the satellite inspiral will be insufficient to change the orientation of the SMBH spin axis as long as the SMBH spin is $> {\\rm few} \\times 0.1$. It follows that the perpendicular orientation of the spin axis has been set by the physical processes before the merger with the satellite, and most likely by the accretion of gas from the Galactic disc. ",
        "conclusions": "A range of theoretical arguments and observational evidence could indicate a satellite infall event within our GC which triggered a brief epoch of strong star formation and AGN activity millions of years ago.  When coupling the newest data -- on the Fermi bubble and the dearth of late-type stars -- to the well-established features of the GC such as the cuspy early-type stellar population, a timeline of the recent dynamical events in the galactic centre emerges.  While the case for a merger of the Milky Way with a satellite galaxy is not beyond reproach, it is a plausible explanation that naturally accounts for both the late- and early-type stellar distributions and the recent violent past of Sgr A*.  This event may not be unique in the evolution of the Milky Way; indeed N-body simulations of the growth of Milky Way-mass galaxies suggest that the present epoch is rife with mergers of relic satellite galaxies with the galactic centre, occurring at a rate of one per few Gyr \\citep{diemand07,shb11}. This implies that there may have been other bursts of hypervelocity star ejections, which can seed a population of ``intragroup stars'' farther out in the halo of the Galaxy. Interestingly, we see tentative evidence in the SDSS archive for a potential set of very late M giants at $\\sim 300\\,{\\rm kpc}$, outside the virial radius of our galaxy~\\citep{palladinoinprep}. Although a followup observation is needed to ensure that these intragroup candidates are not L dwarfs, if these do prove to be very distant giants, they may be provide supporting evidence of a previous minor-merger induced burst of ejected stars $\\sim 10^8$ years ago.  Along similar lines, if satellite infall induced activity is common, then there may be a subset of spiral galaxies which exhibits the signs of the recent onset of the accretion-powered jets. While the longer term X- and $\\gamma$-ray signatures of jets expanding into the intragalactic medium may be too faint to observe in galaxies other than the Milky Way, relatively bright and short lived radio-jets \\citep[$\\sim 0.1\\,{\\rm Myr}$;][]{guo11} may be present in a fraction of up to $\\sim 10^{-4}$ Milky-Way-like spirals, assuming the minor merger rate cited above. Some of these galaxies may be observed serendipitously, during the transient phase associated with the onset of a powerful jet, similar to the case of the previously inactive galaxy J164449.3+573451 that was recently detected by the Swift observatory as a powerful source of beamed emission \\citep{burrows11}. If it can be shown that such a sequence of events occurred in the not so distant past in our Galaxy, it would forever change the paradigm of the Milky Way as an inactive galaxy with an underluminous central SMBH."
    },
    "1107/1107.0117_arXiv.txt": {
        "abstract": "{We report here on the \\XMM\\ observations of the three supergiant fast X-ray transients (SFXT) XTE\\,J1739-302, IGR\\,J08408-4503, and IGR\\,J18410-0535. For the latter source we only discuss some preliminary results of our data analysis. Some interpretation is provided for the timing and spectral behavior of the three sources in terms of the different theoretical models proposed so far to interpret the behavior of the SFXTs.}  \\FullConference{8th INTEGRAL Workshop \"The Restless Gamma-ray Universe\" - Integral 2010,\\\\ September 27-30, 2010\\\\ Dublin Ireland}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  Supergiant fast X-ray transients are a subclass of supergiant X-ray binaries (SgHMXBs) which gained in the past few years a great interest due to their peculiar behavior in the X-ray domain \\citep[see, e.g.,][]{walter07}. The few-hours long outbursts displayed by these sources have been so far attributed to the presence of a NS accreting matter from the extremely clumpy wind of its supergiant companion \\citep[these clumps are expected to be a factor 10$^4$-10$^5$ denser than the homogeneous stellar wind; see e.g.,][]{walter07}. Numerical simulations of supergiant star winds indicate that such high density clumps might results from instabilities in the wind \\citep{oskinova07}. \\citet{bozzo08} proposed that the X-ray variability of the SFXT sources might be driven by centrifugal and/or magnetic ``gating'' mechanisms, that can halt most of the accretion flow during quiescence, and only occasionally permit direct accretion onto the NS. The properties of these gating mechanisms depend mainly on the NS magnetic field and spin period. At odds with the extremely clumpy wind model, in the gating scenario a transition from the regime in which the accretion is (mostly) inhibited to that in which virtually all the captured wind material accretes onto the NS requires comparatively small variations in the stellar wind velocity and/or density, and can easily give rise to a very large dynamic range in X-ray luminosity. Yet another model was proposed for IGR\\,J11215-5952, so far the only SFXT which displays regular periodic outbursts; this model envisages that the outbursts take place when the NS in its orbit crosses a high density equatorial region in the supergiant's wind \\citep{sidoli07}.  The \\XMM\\ observations that we report on this paper were carried out in order to study the quiescent emission of the SFXT sources and gain further insight on the mechanism that drives their peculiar X-ray activity. ",
        "conclusions": "\\label{sec:discussion}  We have presented four deep pointed \\XMM\\ observations of the SFXTs \\xte,\\  \\igrB\\ and \\ax.\\  The first two sources exhibited a very similar behavior. Our spectral and timing analysis revealed that both sources were characterized by a pronounced variability on time-scales of the order of few thousands of seconds. During the \\XMM\\ observation \\xte\\ and \\igrB\\ displayed a number of low-intensity flares, taking place sporadically from a lower persistent emission level. The typical duration of these flares is a few thousands of seconds, and their X-ray flux is a factor of $\\sim$10-30 higher than the persistent quiescent flux. Taking into account the fluxes measured during the brightest outbursts (See Sect.~\\ref{sec:intro}), the total dynamic range of these two sources is $>$10$^4$.  The hardness intensity diagrams of \\xte\\ and \\igrB,\\ together with the results of the rate resolved analysis carried out in Sect.~\\ref{sec:data}, showed that the variations in the X-ray flux measured during the XMM-Newton observations were accompanied in both cases by a change in the spectral properties of the sources. In particular we noticed that the source hardness ratios were increasing significantly with the X-ray flux. A fit to the rate-resolved spectra with a CUTOFFPL model revealed that this behavior originates from a change in the CUTOFFPL photon index, $\\Gamma$, rather than in a variation of the absorption column density. These results indicate that the timing and spectral variability of the two sources during the quiescence is qualitatively similar to that observed during the outbursts (see Sect.~\\ref{sec:intro}). Since most of the X-ray emission from these sources is produced by the accretion of matter onto the NS \\citep{bozzo10}, then the accretion process in these sources takes place over more than 4 orders of magnitude in the X-ray luminosity.  The properties of the soft component dominating the X-ray spectrum of \\igrB\\ and \\xte\\ below $\\simeq$2~keV could be reasonably well constrained in the former source, where the absorption column density was relatively low ($<$10$^{22}$~cm$^{\u22122}$), whereas in the case of \\xte\\ the detection of this component was less significant. However, the similarity in the timing and spectral behavior observed in the quiescent state of the two sources argued in favor of adopting the same spectral model for both of them. We suggested that a MKL component provide a reasonable description of the data, as it would represent the contribution to the total X-ray emission of the shocks in the wind of the supergiant companion. The results of the fits with this model to the data of the three observations inferred a temperature of the MKL component and an emitting region comparable with the values found also in the case of the SFXT AX\\,J1845.0-0433 \\citep{zurita09}. Similar soft spectral components have been detected in many other HMXBs and SGXBs. In a few cases, the detection of a number of prominent emission lines in the high resolution X-ray spectra of these sources carried out with the gratings onboard \\chan\\ and the RGS onboard \\XMM\\ \\citep[see e.g.,][]{watanabe06} have convincingly demonstrated that these components are produced by the stellar wind around the NS, and proved to be a powerful diagnostic to probe the structure and composition of the stellar wind in these systems. The statistics of the present XMM-Newton observations is far too low to permit a similar in-depth study of the stellar wind in the case of \\xte\\ and \\igrB,\\ and further longer X-ray observations of these sources are probably required in order to firmly establish the nature of the soft spectral components in the quiescent emission of the SFXT sources.  \\ax\\ underwent a bright X-ray flare during the \\XMM\\ observation, and we reported here some preliminary results of the analysis of these data. The time-resolved spectral analysis reported in Sect.~\\ref{sec:axresults} showed that during this event the absorption column density of the X-ray emission from the source displayed dramatic changes, whereas the PL photon index remained virtually constant up to $T$$\\simeq$19000~s. These results suggest that the event might have been caused by the accretion of a massive clump, similar to the case of IGR\\,J17544-2619 \\citep{rampy09}. A major change in the spectral properties of the source occurred during the latest 20~ks of observation, in which our analysis revealed an abrupt flattening of the PL photon index and the emergence of a prominent iron-line at 6.4~keV. This spectral transition might indicate that at the end of the flare \\ax\\ underwent some kind of obscuration event, similar to what was observed previously in the case of IGR\\,J16479-4514 \\citep{bozzo08b}. A more refined \\XMM\\ data analysis is currently ongoing in order to understand deeply the origin of the flare recorded from \\ax.\\"
    },
    "1107/1107.2112_arXiv.txt": {
        "abstract": "The symmetron is a scalar field associated with the dark sector whose coupling to matter depends on the ambient matter density. The symmetron is decoupled and screened in regions of high density, thereby satisfying local constraints from tests of gravity, but couples with gravitational strength in regions of low density, such as the cosmos. In this paper we derive the cosmological expansion history in the presence of a symmetron field, tracking the evolution through the inflationary, radiation- and matter-dominated epochs, using a combination of analytical approximations and numerical integration. For a broad range of initial conditions at the onset of inflation, the scalar field reaches its symmetry-breaking vacuum by the present epoch, as assumed in the local analysis of spherically-symmetric solutions and tests of gravity. For the simplest form of the potential, the energy scale is too small for the symmetron to act as dark energy, hence we must add a cosmological constant to drive late-time cosmic acceleration. We briefly discuss a class of generalized, non-renormalizable potentials that can have a greater impact on the late-time cosmology, though cosmic acceleration requires a delicate tuning of parameters in this case. ",
        "introduction": "The standard $\\Lambda$CDM cosmology explains cosmic acceleration with a cosmological constant, a logically sound but seemingly highly tuned explanation.  Therefore there has been much interest in exploring the possibility that the cosmic acceleration could be caused by a previously unobserved dynamical component of the universe~\\cite{Ratra:1987rm}$-$\\cite{Das:2005yj}. At the same time, nearly massless gravitationally coupled scalars are generically predicted to exist by many theories of high energy physics.  No experimental sign of such a fundamental scalar has yet been seen, in spite of tests designed to detect solar system effects and fifth forces that would generally be expected if such scalars were present~\\cite{Fischbach:1999bc,Will:2005va}.  If light scalars exist, they must utilize a screening mechanism~\\cite{Khoury:2010xi} to hide themselves from local experiments. Screening mechanisms rely on non-linearities whose behavior depends on the ambient matter density. In regions of high density, such as the local environment, the scalars develop non-linearities that effectively decouple them from matter. In regions of low density, such as the cosmos, the scalars couple to matter with gravitational strength and mediate a long-range force, thereby affecting the nature of gravity and the growth of structure on large scales~\\cite{Jain:2010ka}.  Aside from cosmology, these theories find independent motivation in the vast experimental effort aimed at testing the fundamental nature of gravity at long wavelengths~\\cite{Will:2005va}. Viable screening theories make novel predictions for local gravitational experiments. The subtle nature of these signals have forced experimentalists to rethink the implications of their data~\\cite{eotvos,Upadhye:2006vi} and have inspired the design of novel experimental~\\cite{Gies:2007su}$-$\\cite{admx} and observational tests~\\cite{claire}$-$\\cite{Jain:2011ji}. The theories of interest offer a rich spectrum of testable predictions, from laboratory to extra-galactic scales.  Only a handful of successful and robust screening mechanisms have been proposed to date~\\cite{Khoury:2010xi}. The first is the Vainshtein mechanism~\\cite{vainshtein,ags,ddgv}, which works when the scalars have derivative self-couplings which become important in the vicinity of massive sources such as the Earth. The strong coupling boosts the kinetic terms, so that after canonical normalization the coupling of fluctuations to matter is weakened.  This mechanism is central to the phenomenological viability of brane-world modifications of gravity~\\cite{Dvali:2000hr}$-$\\cite{Agarwal:2011mg}, massive gravity~\\cite{ddgv},~\\cite{Babichev:2009us}$-$\\cite{Koyama:2011xz} (see \\cite{Hinterbichler:2011tt} for a review), degravitation models~\\cite{Dvali:2002pe}$-$\\cite{Patil:2008sp} and galileon theories~\\cite{galileon}$-$\\cite{Goon:2011qf}. See~\\cite{Afshordi:2008rd}$-$\\cite{Brax:2011sv} for some phenomenological implications.  A second screening mechanism is the chameleon mechanism~\\cite{Khoury:2003aq}$-$\\cite{Brax:2010gi}, which works when the scalars are non-minimally coupled to matter in such a way that their effective mass depends on the local matter density. Deep in space, where the local mass density is low, the scalars are light and display their effects, but near the Earth, where experiments are performed, and where the local mass density is high, they acquire a mass, making their effects short range and unobservable. The effective coarse-grained description of chameleon theories, including careful considerations of averaging, has been derived in~\\cite{Mota:2006fz}. Chameleonic vector fields, such as gauged $B-L$, have been proposed in~\\cite{nelsonchamvec}. See~\\cite{Hinterbichler:2010wu} for an attempt at realizing the chameleon in string compactifications.  Recently, two of us have proposed a third screening mechanism, called the symmetron~\\cite{Hinterbichler:2010es}, based in part on earlier work~\\cite{Olive:2007aj,Pietroni:2005pv}. For this  mechanism to operate, the vacuum expectation value (VEV) of the scalar must depend on the local mass density. The VEV becomes large in regions of low mass density, and small in regions of high mass density. In addition, the coupling of the scalar to matter is proportional to the VEV, so that the scalar couples with gravitational strength in regions of low density, but is decoupled and screened in regions of high density.  This is achieved through the interplay of a symmetry-breaking potential~\\cite{Hinterbichler:2010es}, \\be V(\\phi) = -\\frac{1}{2}\\mu^2\\phi^2 + \\frac{1}{4}\\lambda\\phi^4\\,, \\label{quarticpot} \\ee and a $\\mathds{Z}_2$-invariant universal conformal coupling to the trace of the matter stress tensor, $\\sim\\phi^2 T^\\mu_{\\;\\mu}/M^2$, so a local matter density contributes to the effective mass of the scalar. In vacuum, the scalar acquires a VEV $|\\phi| = \\phi_0 \\equiv \\mu/\\sqrt{\\lambda}$, which spontaneously breaks the reflection symmetry of the lagrangian $\\phi\\rightarrow -\\phi$.  In the presence of sufficiently high ambient density, on the other hand, the potential does not break the symmetry and the scalar is trapped near $\\phi =0$. This is shown in Figure~\\ref{sympot}. In addition, $\\delta\\phi$ fluctuations couple to matter as $({\\phi_{\\rm VEV}}/M^2)\\delta\\phi\\; \\rho$. Hence, symmetron perturbations are weakly coupled in high density backgrounds and relatively more strongly coupled in low density backgrounds.  The symmetron naturally takes the form of an effective field theory. The potential comprises the most general renormalizable terms invariant under the $\\mathds{Z}_2$ symmetry $\\phi\\rightarrow -\\phi$. The coupling to matter is the leading such coupling compatible with the symmetry.  It is non-renormalizable, suppressed by the mass scale $M$, thus the symmetron is an effective theory with cutoff $M$. Remarkably, tests of gravity constrain this cutoff to be around the GUT scale~\\cite{Hinterbichler:2010es}, so any GUT theory with a low energy scalar might be expected to yield a symmetron-type lagrangian at low energies. From this point of view, symmetron theories are more natural-looking than chameleon models. As with all non-supersymmetric scalars, however, the coupling to matter generates large quantum corrections to the mass which must be fine-tuned away.  Symmetron theories predict a host of observational predictions, some of which are distinguishable from other screening mechanisms. Although the local environment is dense, and therefore one in which the symmetron-mediated force is weak, the symmetron nevertheless leads to small deviations from General Relativity in the solar system. As reviewed in Section~\\ref{tests}, the predicted signals for Lunar Laser Ranging and Mercury's perihelion precession are just below current bounds and within reach of next-generation experiments~\\cite{Hinterbichler:2010es}. On the other hand, the signal from binary pulsars is much weaker and should not be observed, since both neutron stars are screened. This makes the symmetron distinguishable from over-the-counter scalar-tensor theories, such as Brans-Dicke (BD) theories, for which solar system and binary pulsar signals are comparable. The symmetron solar system predictions are also distinguishable from chameleon and Vainshtein screening~\\cite{Hinterbichler:2010es}.  Symmetrons, like chameleons, predict effective macroscopic violations of the equivalence principle between large (screened) and small (unscreened) bodies. Dwarf galaxies in low-density environments, in particular, offer a fertile playground for testing these ideas~\\cite{Hui:2009kc,Jain:2011ji}. The stars are oblivious to the symmetron whereas the hydrogen gas experiences an additional force, resulting in an enhanced rotational velocity for the gas~\\cite{Hui:2009kc}. Meanwhile, the infall motion of a dwarf galaxy can lead to a segregation of the stellar disk from the dark matter and the hydrogen gas~\\cite{Jain:2011ji}.  \\begin{figure}[t] \\centering \\subfloat[Effective potential in regions of high ambient density. The VEV is 0.]{\\includegraphics[scale=.8]{symmetric_potential.pdf}} \\qquad \\subfloat[Effective potential in regions of low ambient density.]{\\includegraphics[scale=.8]{symmetry_broken_potential.pdf}} \\caption{Schematic plots of the symmetron effective potential, illustrating the symmetry breaking phase transition.} \\label{sympot} \\end{figure}  We begin with an extensive overview of symmetron physics in Section~\\ref{symrev}, reviewing the spherically-symmetric solution (Section~\\ref{sphsol}) and constraints from solar system tests of gravity (Section~\\ref{tests}). Section~\\ref{jordan}, in particular, offers a novel perspective of the symmetron in terms of Jordan-frame variables. In this frame, the action describes a Brans-Dicke type scalar-tensor theory, with a field-dependent Brans-Dicke parameter. The screening mechanism is understood as the Brans-Dicke parameter becoming large in regions of high-density, corresponding to the decoupling of the scalar. In this respect, the symmetron is qualitatively similar to the Vainshtein mechanism, though its non-linearities stem from potentials rather than derivative self-couplings.  The remainder of the paper focuses on the cosmological evolution in the presence of a symmetron field. Specifically, we are interested in checking whether a symmetron leads to an acceptable cosmology and, more interestingly, whether it can act as dark energy driving cosmic acceleration. We derive the evolution of the scalar field from the inflationary epoch until the present. The dependence of the symmetron effective potential on the matter density has implications for the cosmological evolution of the scalar field. Since the matter density redshifts in time, the effective potential is time-dependent and results in a phase transition when the matter density falls below a critical value. We will choose parameters such that the phase transition occurs in the recent past ($z_{\\rm tran} \\,\\lsim \\, 1$).  A key question is whether the evolution allows the scalar field to reach the symmetry-breaking vacuum by the present epoch, as assumed in the analysis of solar system tests and other phenomenological studies. The short answer is that the coupling to matter efficiently drives the field towards the symmetry-restoring point, so that for a broad range of initial conditions it reaches $\\phi\\simeq 0$ well before the phase transition. How the field makes it there, however, is an interesting story, as summarized below.  In Section~\\ref{cosmoevol} we describe the symmetron evolution during the radiation- and matter-dominated eras of standard big bang cosmology.  The effective mass squared of the symmetron due to its coupling to matter fields is $m_{\\rm eff}^2 \\sim T^\\mu_{\\;\\mu}/M^2 \\sim (1-3w)\\rho/M^2$, where $w$ is the equation of state. During the radiation-dominated era, $T^\\mu_{\\;\\mu}$ is dominated by the (subdominant) non-relativistic component, hence $m_{\\rm eff}^2 \\sim \\Omega_{\\rm m} H^2M_{\\rm Pl}^2/M^2$. At early times, therefore, $m_{\\rm eff}^2 \\ll H^2$, and the symmetron remains essentially frozen at its initial value, denoted by $\\phi_{\\rm rad-i}$. As the universe cools and $\\Omega_{\\rm m}$ increases, eventually $m_{\\rm eff}^2 \\simeq H^2$, and the symmetron starts rolling and undergoes damped oscillations around $\\phi = 0$. Since $M \\ll M_{\\rm Pl}$, this rolling phase is triggered deep in the radiation-dominated epoch, prior to Big Bang Nucleosynthesis (BBN). From this moment onwards, the field amplitude decreases by a factor of $(M/M_{\\rm Pl})^{3/2}$ until matter-radiation equality. The damped oscillations continue during the matter-dominated era, and the field amplitude decreases by an additional factor of $10^{-3}$ until the present epoch. Thus, for a broad range of initial field values, the symmetron ends up close to the origin by the onset of the phase transition.  In Section~\\ref{phasetran}, we follow the evolution of the symmetron through the phase transition. For the quartic potential~(\\ref{quarticpot}), the symmetron nicely tracks the effective minimum and settles to the symmetry-breaking vacuum by the present time. However, the energy scale of the potential is too small to drive cosmic acceleration, and the symmetron backreaction is negligible. Hence we must add a suitable cosmological constant, thereby making the expansion history indistinguishable from $\\Lambda$CDM cosmology.  We also consider generalized, non-renormalizable forms of the potential, with scales chosen to have a greater impact on late-time cosmology: $i)$ the potential energy difference between maximum and minima is of order $H_0^2M_{\\rm Pl}^2$, as desired to act as dark energy; and $ii)$ the mass of small fluctuations around the minima is of order $H_0$, such that the symmetron can impact the growth of structure on the largest observable scales. The resulting potential is sharply peaked at $\\phi = 0$, and displays very shallow minima. Because of this asymmetric form, shortly after the onset of the phase transition the symmetron overshoots the minimum and reaches a value of order $M_{\\rm Pl}$. As a result, the field does not converge to the symmetry-breaking vacuum by the present time, which is problematic for tests of gravity. Meanwhile, the symmetron energy density quickly becomes irrelevant after the transition. One way to circumvent these problems is by requiring the phase transition to occur in the very recent past ($z_{\\rm tran} \\ll 1$), but this requires a delicate choice of parameters.  In Section~\\ref{inf} we go back to the early universe and describe the evolution of the symmetron during a period of inflation and subsequent reheating. Since the effective mass $m_{\\rm eff}^2 \\simeq 12 H^2_{\\rm inf}M_{\\rm Pl}^2/M^2$ is $\\gg H^2_{\\rm inf}$ during inflation, the scalar undergoes rapid oscillations with exponentially decaying amplitude, thanks to Hubble damping. The end result is that a field initially displaced from $\\phi = 0$ ends up exponentially close to this point. At the end of inflation, however, as the inflaton oscillates around its minimum, its equation of state also oscillates between $w = -1$ and $w = 1$. In turn, $m_{\\rm eff}^2$ oscillates between negative and positive values, which can result in amplification.  With a simple parametrization of $w(t)$, we find that the symmetron evolution can be described at early times by a Mathieu equation, with solution given by the Mathieu cosine function. As usual with Mathieu functions, the stability behavior depends on parameter choices, in this case  the inflation mass (setting the rate of oscillations in $m_{\\rm eff}^2$), its decay rate (setting the duration of the unstable phase) and the effective symmetron mass during inflation. We place a constraint on these parameters by demanding that the symmetron displacement be within effective field theory bounds.  We summarize our results and discuss future directions in Section \\ref{conclsec}. ",
        "conclusions": ""
    },
    "1107/1107.0926_arXiv.txt": {
        "abstract": "\\noindent We study supersymmetric extensions of the Standard Model with small R-parity and lepton number violating couplings which are naturally consistent with primordial nucleosynthesis, thermal leptogenesis and gravitino dark matter. We consider supergravity models where the gravitino is the lightest superparticle followed by a bino-like next-to-lightest superparticle (NLSP). Extending previous work we investigate in detail the sensitivity of LHC experiments to the R-parity breaking parameter~$\\zeta$ for various gluino and squark masses. We perform a simulation of signal and background events for the generic detector \\software{DELPHES} for which we implement the finite NLSP decay length. We find that for gluino and squark masses accessible at the LHC, values of $\\zeta$ can be probed which are one to two orders of magnitude smaller than the present upper bound obtained from astrophysics and cosmology. ",
        "introduction": "Supersymmetric extensions of the Standard Model with broken R-parity have a rich phenomenology \\cite{Hall:1983id,Ross:1984yg,Ellis:1984gi,Allanach:2003eb, Barbier:2004ez}. In most models rather large R-parity violating couplings are considered, which lead to prompt decays of the lightest superparticle in the detector. In models where small R-parity violating interactions generate neutrino masses, macroscopic decay lengths up to \\unit[1]{mm} are obtained \\cite{deCampos:2007bn}. In the case of gauge mediated supersymmetry breaking, R-parity violating decays then compete with R-parity conserving decays where the final state contains a gravitino \\cite{Hirsch:2005ag}.  If the gravitino is the lightest superparticle, its decays are doubly suppressed by the Planck mass and the small R-parity breaking parameter. Hence, its lifetime typically exceeds the age of the universe by many orders of magnitude, so that it remains a viable dark matter candidate \\cite{Takayama:2000uz}. In the case of very small R-parity breaking couplings, as they occur if R-parity is spontaneously broken at the grand unification scale, one obtains a consistent cosmology including primordial nucleosynthesis, thermal leptogenesis and gravitino dark matter \\cite{Buchmuller:2007ui}. In the following we shall study the collider phenomenology of these models extending our previous work \\cite{Bobrovskyi:2010ps}.  Decaying gravitino dark matter leads to a diffuse gamma-ray flux \\cite{Takayama:2000uz,Buchmuller:2007ui,Lola:2007rw}. For a bino-like NLSP the matrix elements for gravitino decay and NLSP decay are directly related \\cite{Bobrovskyi:2010ps}. Together with the lower bound on the gravitino lifetime, and the corresponding upper bound on the R-parity breaking parameter $\\zeta$, which is derived from the diffuse gamma-ray flux observed by the Fermi-LAT collaboration \\cite{Abdo:2010nc,Abdo:2010nz}, one obtains a lower bound on the NLSP decay length. The estimates in \\cite{Bobrovskyi:2010ps} led to $\\zeta \\lesssim 3\\times 10^{-8}$, and in the more detailed analysis of gamma-ray lines in \\cite{Vertongen:2011mu} a slightly stronger bound was found, $\\zeta \\lesssim 2\\times 10^{-8}$. For gravitino masses $\\mathcal{O}(\\unit[100]{GeV})$ the upper bound becomes more stringent: $\\zeta \\sim 10^{-9}$ \\cite{Bobrovskyi:2010ps}. This range in the parameter $\\zeta$ corresponds to NLSP decay lengths varying from $\\mathcal{O}(\\unit[50]{cm})$ to $\\mathcal{O}(\\unit[500]{m})$.  Large macroscopic decay lengths are of great help in the search for decaying NLSPs. This remains true if the decay length is larger than the size of the detector since a sizeable fraction of NLSPs may still decay inside the detector. This has been studied for neutral \\cite{Ishiwata:2008tp} as well as charged \\cite{Asai:2009ka} NLSPs. Neutralino decay lengths varying from \\unit[0.1]{mm} to \\unit[100]{m} also arise in models with generalized gauge mediated supersymmetry breaking \\cite{Meade:2010ji}. Alternatively, charged \\cite{Desch:2010gi} and neutral \\cite{Bomark:2011ye} NLSP decays have been studied for models where the decay lengths are so small that no displaced vertices are observed and R-parity breaking Yukawa couplings determine the hierarchy of decay channels. In this case multi-lepton events, and their flavour structure, are of crucial importance.  The subject of this paper is a quantitative analysis of neutralino NLSP decays at the LHC in the case of very small R-parity breaking, $\\zeta \\lesssim 3\\times 10^{-8}$. The goal is the determination of the sensitivity in $\\zeta$ for varying gluino and squark masses. To achieve this we perform a simulation for the generic detector \\software{DELPHES} \\cite{Ovyn:2009tx} of signal and background events. We focus on events with a clean signature: cascade processes with jets where one of the produced neutralino NLSPs decays into $Z$-boson and neutrino, with the subsequent decay of the $Z$-boson into a muon pair. This allows us to determine a conservative 5$\\sigma$ discovery range. Finally, we estimate the discovery reach of the LHC if all NLSP decays are taken into account.  The paper is organized as follows. In Section~\\ref{sec:r_parity_breaking} we recall the main ingredients of the considered model of bilinear R-parity breaking, as well as the dominant NLSP decays into gauge bosons and leptons. In Section~\\ref{sec:signatures_and_reach} we discuss qualitative signatures of the events: NLSP $\\beta\\gamma$ distribution, $\\slashed p_T$ spectrum, and the number of final state leptons. In Section~\\ref{sec:signal_and_background}, five representative benchmark points are defined and the simulation of signal and background events is described with emphasis on the muon reconstruction. The numerical results for the chosen NLSP decay channel $Z(\\mu^+\\mu^-)\\nu$ are given in Section~\\ref{sec:neutrsearch}, together with estimates for the discovery reach if all NLSP decays are included. We conclude with a summary in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have studied the decays of quasi-stable neutralinos in supersymmetric extensions of the Standard Model with small R-parity and lepton number violating couplings which are consistent with primordial nucleosynthesis, thermal leptogenesis and gravitino dark matter. As representative examples with neutralino NLSP we have considered five benchmark points, HH27--HH80, with gluino and squark masses ranging from $\\unit[650]{GeV}$ to $\\unit[1800]{GeV}$.  For the considered benchmark points the neutralino is bino-like, and the R-parity breaking parameter $\\zeta$, which governs the neutralino decays $\\chi_1^0 \\rightarrow W^\\pm l^\\mp$ and $\\chi_1^0 \\rightarrow Z \\nu$, is directly related to the partial gravitino decay width into photon and neutrino.  The main goal of the present work has been to determine the range of the parameter $\\zeta$ which can be probed at the LHC, for varying superparticle masses. As a conservative starting point, we have focused on events with a clean signature: cascade processes with jets where one of the produced neutralino NLSPs decays into $Z$-boson and neutrino, with a subsequent decay of the $Z$-boson into a muon pair.  In Section~\\ref{sec:signatures_and_reach} we have studied the qualitative signatures of these events, the $\\beta\\gamma$ distribution of the produced NLSPs, the $\\slashed p_T$-spectrum and the number of leptons in the final state. A detailed simulation of signal and background events for the generic detector \\software{DELPHES} (with CMS tune) has been described in Section~\\ref{sec:signal_and_background}, with emphasis on the reconstruction of muons. The crucial new element of our analysis has been the implementation of the finite NLSP decay length.  The numerical results of our simulations have been summarized in Section~\\ref{sec:neutrsearch}. It is very instructive to look at the spacial distribution of decay vertices which have been given for different benchmark points and different values of $\\zeta$. One can clearly see that the sensitivity extends from decay lengths $\\mathcal{O}(\\unit[50]{cm})$, with NLSP decays mostly inside the detector, to values $\\mathcal{O}(\\unit[500]{m})$ where almost all NLSPs decay outside the detector.  The results for the discovery reach for quasi-stable neutralino NLSPs roughly agree with the simple estimates which one obtains from the branching ratios into the $Z(\\mu^+\\mu^-)\\nu$ final state together with the assumption that these events are background free. It is remarkable that already with $\\unit[10]{fb^{-1}}$ a $5\\sigma$ discovery is possible for squark and gluino masses of \\unit[830]{GeV} and an R-parity breaking parameter $\\zeta = 3\\times10^{-9}$, which is one order of magnitude smaller than the present astrophysical bound.  It is likely that the severe restriction of our simulation to $Z(\\mu^+\\mu^-)\\nu$ final states can be relaxed and that a much larger fraction of events can be used for the analysis. Assuming optimistically that 10 decays inside the detector are sufficient for a discovery, one would be sensitive to values of $\\zeta$ down to $3\\times10^{-10}$ for squark and gluino masses of \\unit[830]{GeV} and $\\unit[10]{fb^{-1}}$. Alternatively, for gluino and squark masses of \\unit[1480]{GeV} one could probe values of the R-parity breaking parameter down to $3\\times10^{-9}$. For the same quark masses a luminosity of $\\unit[100]{fb^{-1}}$ would improve the sensitivity in $\\zeta$ by a factor of three. We conclude that for gluino and squark masses accessible at the LHC values of the R-parity breaking parameter can be probed which are far below the present upper bounds obtained from astrophysics and cosmology.  \\subsection*"
    },
    "1107/1107.0860_arXiv.txt": {
        "abstract": "We present a model of optically thin, two-temperature, accretion flows using an exact Monte Carlo treatment of global Comptonization, with seed photons from synchrotron and bremsstrahlung emission, as well as with a fully general relativistic description of both the radiative and hydrodynamic processes. We consider accretion rates for which the luminosities of the flows are between $\\sim 10^{-3}$ and $10^{-2}$ of the Eddington luminosity. The black hole spin parameter strongly affects the flow structure within the innermost $\\simeq 10$ gravitational radii. The resulting large difference between the Coulomb heating in models with a non-rotating and a rapidly rotating black hole is, however, outweighed by a strong contribution of compression work, much less dependent on spin.  The consequent reduction of effects related to the value of the black spin is more significant at smaller accretion rates. For a non-rotating black hole, the compressive heating of electrons dominates over their Coulomb heating, and results in an approximately constant radiative efficiency of $\\approx 0.4$ per cent in the considered range of luminosities. For a rapidly rotating black hole, the Coulomb heating dominates, the radiative efficiency is $\\simeq 1$ per cent and it slightly increases (but less significantly than estimated in some previous works) with increasing accretion rate. Our study neglects the direct heating of electrons, which effect can lead to larger differences between the radiative properties of models with a non-rotating and a rapidly rotating black hole than estimated here. Flows with the considered parameters produce rather hard spectra, with the photon spectral index $\\Gamma \\sim 1.6$, and with high energy cut-offs at several hundred keV. We find an agreement between our model, in which the synchrotron emission is the main source of seed photons, and observations of black-hole binaries in their hard states and AGNs at low luminosities. In particular, our model predicts a hardening of the X-ray spectrum with increasing luminosity, as indeed observed below $\\sim 0.01 \\ledd$ or so in both black-hole binaries and AGNs. Also, our model approximately reproduces the luminosity and the slope of the X-ray emission in Cen A. ",
        "introduction": "\\label{intro}  Optically thin, two-temperature accretion flows have been considered as an explanation of a variety of black hole systems and a substantial work has been done for investigation of their dynamical and spectral properties (see, e.g., Narayan \\& McClintock 2008 for a review). However, the developed models still involve several approximations which significantly reduce their accuracy. These involve the use of a pseudo-Newtonian potential of Paczy\\'nski \\& Wiita (1980), which fails in the innermost region (particularly if rotation of the black hole is considered), where most of the gravitational energy is dissipated, as well as local approximations of Comptonization, which appear to be particularly incorrect in optically thin flows (see Xie et al.\\ 2010; hereafter X10). Previous attempts to improve these two weaknesses are discussed in Section \\ref{previous}.  In this paper, we extend our previous treatment of global Comptonization from X10 by using the hydrodynamical model from Manmoto (2000; hereafter M00), and develop a self-consistent model involving a fully general-relativistic (GR) description of both the hydrodynamical and the radiative processes. We consider moderate values of accretion rate, at which Coulomb coupling between ions and electrons is relatively weak compared to the viscous heating of ions and, therefore, the ion temperature (and hence the total pressure) is not affected significantly by the details of the description of radiative processes. This allows for a slightly simplified treatment of the flow structure in our computations leading to a self-consistent solution.  Already the foundational papers proposed that tenuous, two-temperature flows may be responsible for the hard spectral states of black-hole binaries (Ichimaru 1977; see also Narayan \\& Yi 1995) and for the low nuclear luminosities in radio galaxies with large radio lobes (Rees et al.\\ 1982). We consider two ranges of the key parameters (black hole mass and accretion rate) which should be relevant for these two major application areas. For each of these two cases, we illustrate the impact of the black hole spin by considering a non-rotating and a maximally-rotating black hole. On the other hand, we neglect here the dependence on some other parameters of hot-flow models.  In particular, we consider only weakly magnetized flows, with the magnetic pressure of 1/10th of the total pressure, and we neglect the direct viscous heating of electrons; see Section \\ref{delta} for a discussion of these two assumptions. Furthermore, we neglect outflows, which may play an important role in some systems (e.g.\\ Yuan, Quataert \\& Narayan 2003). Then, our solutions are strongly dominated by the advection of energy by ions, as in the original formulation of the advection-dominated accretion flow (ADAF) model by, e.g., Narayan \\& Yi (1995).     In this paper, we focus on modelling an innermost part of an accretion flow, namely inside $10^3 R_{\\rm g}$, where $R_{\\rm g} = GM/c^2$ is the gravitational radius and $M$ is the black-hole mass, where the bulk of the observed radiation is produced. Global Comptonization has a net effect of Compton heating outside $\\sim\\! 10^4 R_{\\rm g}$ (e.g.\\ Park \\& Ostriker 2001, Yuan, Xie \\& Ostriker 2009, Yuan \\& Li 2011), which effect is not considered here. ",
        "conclusions": "We have developed a fully GR model with global Compton scattering of hot, tenuous accretion flows. The dominant source of seed photons in our model is the synchrotron process. We assume that an outer, optically-thick, accretion disc is truncated at large radii, so irradiation of the hot flow by it is negligible. We have also neglected direct viscous electron heating and outflows. We have considered the accretion rates of $\\dot m = 0.1$ and 0.5, for which our model gives the bolometric luminosities of $L/\\ledd \\simeq (0.4$--$8)\\times 10^{-3}$.  In this range of $L/\\ledd$, we find our models predict X-ray spectra hardening with increasing $L$, which is in agreement with the correlation observed at $L/\\ledd\\la 0.01$ in both black-hole binaries in the hard state and AGNs. This is also in agreement with the suggestion of Sobolewska et al.\\ (2011) that synchrotron seed photons dominate in this range of $L/\\ledd$, whereas an optically-thick accretion disc overlapping with the hot flow provide dominant source of seed photons at higher $L$.  We have compared our models to the low-$L$ FR I radio galaxy Cen A. We have found a good agreement with the observed X-ray spectral slope, but the position of the high-energy break, observed at $\\sim$200 keV, is lower than the model cut-off energy. This may be possibly explained by the hadronic processes, which would be present if the black hole in this source has a high spin.  We have studied the impact of the black hole spin and we find that despite a very strong influence of this parameter on the dynamical properties of the inner part of the flow, the difference between the observed radiation from a flow surrounding a maximally rotating and a non-rotating black hole, with other parameters unchanged, is rather moderate. In particular, their luminosities differ only by a factor of 2--3 (i.e.\\ smaller than a factor of $\\simeq 6$ in Keplerian discs). We find this is due to a strong role of the compressive electron heating, which process only weakly depends on the black-hole spin. We note, however, that a much larger dependence on the spin value is expected in models involving direct heating of electrons, or taking into account hadronic processes.  We stress that a proper treatment of global Compton scattering is crucial for accurate modelling of the flow structure, most importantly, for the electron temperature profile, and for the spectral formation in optically thin flows. In models with $\\dot m = 0.5$, we find the entire inner flow is efficiently cooled by Comptonization when the non-local nature of this process is taken into account. At $\\dot m = 0.1$, global Comptonization is important, but also electron advection effects are strong at $r\\ga 10$, giving rise to large temperature gradients. As compared to the results with the global Comptonization, the local model assuming a slab geometry provides in general a very poor approximation, overestimating the cooling rate by an order of magnitude in the innermost region."
    },
    "1107/1107.0267_arXiv.txt": {
        "abstract": "{The detection of pulsational frequencies in stellar photometry is required as input for asteroseismological modelling. The second short run (SRa02) of the CoRoT mission has provided photometric data of unprecedented quality and time-coverage for a number of O-type stars. } {We analyse the CoRoT data corresponding to three hot O-type stars, describing the properties of their light curves and we search for pulsational frequencies, which we then compare to theoretical model predictions. } {We determine the amplitude spectrum of the data, using the Lomb-Scargle and a multifrequency HMM-like technique. Frequencies are extracted by prewhitening, and their significance is evaluated under the assumption that the light curve is dominated by red noise. We search for harmonics, linear combinations and regular spacings among these frequencies. We use simulations with the same time sampling as the data as a powerful tool to judge the significance of our results. From the theoretical point of view, we use the MAD non-adiabatic pulsation code to determine the expected frequencies of excited modes. } {A substantial number of frequencies is listed, but none can be convincingly identified as being connected to pulsations. The amplitude spectrum is dominated by red noise. Theoretical modelling shows that all three O-type stars can have excited modes but the relation between the theoretical frequencies and the observed spectrum is not obvious. } {The dominant red noise component in the hot O-type stars studied here clearly points to a different origin than the pulsations seen in cooler O stars. The physical cause of this red noise is unclear, but we speculate on the possibility of sub-surface convection, granulation, or stellar wind inhomogeneities being responsible. } ",
        "introduction": "Even though the O-type stars are located in the Hertzsprung-Russell (HR) diagram inside a zone where pulsations are expected, no asteroseismological modelling of these stars could be performed prior to the CoRoT \\citep[Convection, Rotation and planetary Transits,][]{2006ESASP1306...33B, 2009A&A...506..411A} mission. Actually, few examples of variability have been identified. The most obvious ones are the O9.5 V stars \\object{$\\zeta$ Oph} \\citep{1997ApJ...481..406K} and \\object{HD 93521} \\citep{1993A&A...279..148H, 2008A&A...487..659R} and their variability is likely related to non-radial pulsations with periods of a few hours. However, the majority of these detections were made spectroscopically. The reason is that the amplitudes of pulsations are too low to be measured from ground-based photometry. Furthermore, these pulsations can be contaminated by variable stellar winds.  The second short run (SRa02) of the CoRoT satellite in the Asteroseismologic Channel was partly devoted to the investigation of the photometric variability of O-type stars. Pointing towards the anti-centre of the Galaxy, this instrument observed objects belonging to the young open cluster \\object{NGC 2244} inside the \\object{Rosette nebula} and to the surrounding association \\object{Mon OB2}. The CoRoT data of three O-type stars that are part of this run have been analysed in previous papers: \\object{HD 46149} by \\citet{2010A&A...519A..38D}, \\object{HD 47129} by \\citet{2011A&A...525A.101M} and \\object{HD 46202} by \\citet{2011A&A...527A.112B}, the latter paper containing the only forward modelling in terms of seismic data of an O star so far.  In the present paper, we study the three remaining O-type stars that are part of the SRa02 run. One of the targets is \\object{HD~46223}, the hottest member of NGC 2244. Situated at about 1.4--1.7 kpc \\citep{2000A&A...358..553H}, this star was previously quoted in the literature as having a spectral type O5~V \\citep{1956ApJS....2..389H, 1982A&A...107..252B}, O5~((f)) \\citep{1974ApJ...193..113C}, O4~((f)) \\citep{1995ApJ...454..151M} and O4~V((f$^+$)) \\citep{2009A&A...502..937M}. The 9-year spectroscopic campaign by \\citeauthor{2009A&A...502..937M} to investigate multiplicity did not reveal any significant variation of the radial velocity related to the presence of a companion. In addition, Chandra mosaic observations of NGC 2244 \\citep{2008ApJ...675..464W} detect no strong clustering of X-ray sources around HD~46223, leading to the conclusion that there is also no lower-mass neighbour present. According to \\citeauthor{2008ApJ...675..464W}, the small number of close-by X-ray sources could indicate that this star is younger in comparison to the population of the central part of the cluster.  \\begin{table*} \\caption{CoRoT data and results from the analysis for the three stars.} \\label{table corot data} \\centering \\begin{tabular}{lllll} \\hline\\hline & Parameter & HD~46223 & HD~46150 & HD~46966 \\\\ \\hline \\multicolumn{5}{l}{CoRoT data} \\\\ & number of raw data, $N_{\\rm R}$ & 92687 & 92695 & 92688 \\\\ & number of non-flagged data, $N$ & 80403 & 80282 & 82273 \\\\ & duty cycle ($N/N_{\\rm R}$, in \\%) & 86.75 & 86.61 & 88.76 \\\\ & start date non-flagged data (HJD - 2\\,450\\,000)& 4748.488547 & 4748.485591 & 4748.488454 \\\\ & end date non-flagged data (HJD - 2\\,450\\,000) & 4782.819650 & 4782.819657 & 4782.805898 \\\\ & duration, $T$ (d)               & 34.331103   &   34.334066 &   34.317444 \\\\ & freq resolution (d$^{-1}$) & 0.029128    &   0.029126  &   0.029140  \\\\ \\hline \\multicolumn{5}{l}{Spectral window} \\\\ & window peak (\\%) at 2.0 d$^{-1}$    &  3.5  &  3.8  &  3.9 \\\\ & window peak (\\%) at 4.0 d$^{-1}$    &  1.1  &  1.0  &  1.1 \\\\ \\hline \\multicolumn{5}{l}{Stopping criterion: number of frequencies} \\\\ & AIC$_{\\rm c}$ & 860 & 821 & 595 \\\\ & BIC           & 491 & 444 & 277 \\\\ & HQC           & 662 & 625 & 427 \\\\ \\hline \\multicolumn{5}{l}{Number of red-noise significant frequencies} \\\\ &               & 59 & 50 & 53 \\\\ \\hline \\multicolumn{5}{l}{Number of red-noise significant frequencies present in both halves of the observing run} \\\\ &               & 10 & 7 & 3 \\\\ \\hline \\multicolumn{5}{l}{Red noise -- fit parameters} \\\\ & $\\alpha_0$ (counts)    & $442 \\pm 31$  & $430 \\pm 65$  & $416 \\pm 85$  \\\\ & $\\tau$ (d)   & $0.09 \\pm 0.01$ & $0.08 \\pm 0.02$ & $0.17 \\pm 0.05$ \\\\ & $\\gamma$  & $0.96 \\pm 0.02$  & $0.96 \\pm 0.03$ & $0.91 \\pm 0.02$ \\\\ \\hline \\end{tabular} \\end{table*}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[bb=20 260 555 685,clip]{Ostars_Fig1.pdf}} \\caption{Spectral window of the HD~46223 observations. The inset zooms in on the low-frequency region. } \\label{fig spectral window} \\end{figure}    The next target is \\object{HD 46150}, the second hottest star in NGC 2244. Literature values for the spectral type are between O5~V((f)) \\citep{1990ApJ...364..626U, 2004ApJS..151..103M} and O5.5~V((f)) \\citep{1977ApJ...213..438C, 1980ApJ...242.1063G}. The detailed study of \\cite{2009A&A...502..937M} confirms the O5.5~V((f)) spectral type. Radial velocity variations have been detected in this star, but it is not clear if these are due to binarity or motions in the stellar atmosphere \\cite[see][and references therein]{2009A&A...502..937M}. \\citeauthor{2009A&A...502..937M} monitored this star but could not find any significant period in their data. They note the lack of spectral variability over short timescales ($\\sim$\\,7 days).  The third target studied in the present paper is the late O-type star \\object{HD 46966}, belonging to the Mon OB2 association. The spectral classification generally agrees on O8.5~V \\citep{1977ApJ...214..759C, 1980ApJ...242.1063G, 2009A&A...502..937M} or O8~V \\citep{1999A&AS..137..521M, 2004ApJS..151..103M}. Moreover, spectroscopic monitoring reveals no clues on the existence of a secondary component \\citep{2009A&A...502..937M}.  These three stars are among the youngest \\citep[1--6 Myr,][]{2009MNRAS.394.2127B} and most massive observed by the CoRoT satellite. They therefore have the potential to extend our knowledge about pulsations into the higher-mass domain. The CoRoT light curves are of unprecedented quality and have a long observational time span (longer than the rotational cycle). They thus provide the ideal data set to reveal the presence or absence of such pulsations.   We first present, in Sect.~\\ref{sect HD 46223}, the detailed frequency analysis of the HD~46223 CoRoT light curve. In Sect.~\\ref{sect HD 46150} and Sect.~\\ref{sect HD 46966}, we give the frequency analysis of HD~46150 and HD~46966 respectively. Section~\\ref{sect asteroseismological modelling} is devoted to determining the theoretically expected pulsation frequencies. We discuss the results in Sect.~\\ref{sect discussion} and in Sect.~\\ref{sect conclusions} we provide the conclusions of our research. ",
        "conclusions": "\\label{sect conclusions}  The CoRoT light curves of the hot O-type stars HD~46223, HD~46150 and HD~46966 were analysed using standard methods to search for pulsation frequencies. The detection of such frequencies would allow an asteroseismological interpretation. However, the results show that most of the variations of these stars are of a stochastic nature. The only significant exception to the above is the possible rotation period detected for HD~46223 and HD~46966, although these conclusions are not strong and still require confirmation. Evidence for a rotation period in the HD~46150 data is even less strong.  The periodogram of these three stars is clearly different from that of a classical pulsator such as HD~46202. Most of the power in our periodograms is not due to pulsation, but is more correctly described by red noise. Based on the limited number of stars studied so far, a trend is suggested: the earliest O stars show red noise, while the later O-types have pulsational frequencies of the $\\beta$ Cep type. The switch-over occurs around spectral type O8.  The specific physical cause of this noise in O-type stars is at present unclear. We point out the possibilities of sub-surface convection, granulation or inhomogeneities in the wind. All three options are highly speculative and therefore await confirmation by detailed modelling. The CoRoT data of these three hot stars present interesting challenges for our understanding of the outer photosphere and inner wind of early-type stars."
    },
    "1107/1107.1261_arXiv.txt": {
        "abstract": "Using deep Keck spectroscopy of Lyman break galaxies selected from infrared imaging data taken with the Wide Field Camera 3 onboard the Hubble Space Telescope,  we present new evidence for a reversal in the redshift-dependent fraction of star forming galaxies with detectable Lyman alpha (Ly$\\alpha$) emission in the redshift range $6.3 < z <8.8$.  Our earlier surveys with the DEIMOS spectrograph demonstrated a significant increase with redshift in the fraction of line emitting galaxies over the interval $4<z<6$, particularly for intrinsically faint systems which dominate the luminosity density. Using the longer wavelength sensitivities of LRIS and NIRSPEC, we have targeted 19 Lyman break galaxies selected using recent WFC3/IR data whose photometric redshifts are in the range $6.3<z<8.8$ and which span a wide range of intrinsic luminosities. Our spectroscopic exposures typically reach a 5$\\sigma$ sensitivity of $<50$ \\AA\\ for the rest-frame equivalent width (EW) of Ly$\\alpha$ emission. Despite the high fraction of emitters seen only a few hundred million years later, we find only 2 convincing and 1 possible line emitter in our more distant sample. Combining with published data on a further 7 sources obtained using FORS2 on the ESO VLT, and assuming continuity in the trends found at lower redshift, we discuss the significance of this apparent reversal in the redshift-dependent Ly$\\alpha$ fraction in the context of our range in continuum luminosity. Assuming all the targeted sources are at their photometric redshift and our assumptions about the Ly$\\alpha$ EW distribution are correct, we would expect to find so few emitters in less than 1\\% of the realizations drawn from our lower redshift samples. Our new results provide further support for the suggestion that, at the redshifts now being probed spectroscopically, we are entering the era where the intergalactic medium is partially neutral. With the arrival of more sensitive multi-slit infrared spectrographs, the prospects for improving the statistical validity of this result are promising. ",
        "introduction": "\\label{sec:intro}  Determining when neutral hydrogen in the intergalactic medium (IGM) was reionized is an important question in observational cosmology and a precursor to understanding whether star forming galaxies provided the necessary UV photons \\citep{Robertson10}. One of  the most practical probes of reionization with current facilities utilizes the frequency of occurrence of Ly$\\alpha$ emission in star forming galaxies.  As Ly$\\alpha$ photons are resonantly scattered by neutral hydrogen, the abundance of Ly$\\alpha$ emitters should decrease as observations probe into the era where there is neutral gas (e.g., \\citealt{Malhotra04}, \\citealt{Ouchi10}, \\citealt{Kashikawa11}).  The recent discovery of large numbers of candidate galaxies beyond $z\\simeq7$ through multi-color imaging undertaken with the infrared Wide Field Camera (WFC3/IR) onboard Hubble Space Telescope (HST) (e.g., \\citealt{Bouwens10b,Bunker10,Finkelstein10,McLure10}) now makes it feasible to track the occurrence of Ly$\\alpha$ line emission to interesting redshifts where neutral hydrogen may be present.  Of course, astrophysical factors other than a neutral IGM can also affect the presence of Ly$\\alpha$ emission. Because of this, an alternative approach for gauging when reionization occurred, introduced by \\citet{Stark10} (hereafter Paper I), is to spectroscopically measure {\\it the fraction of Ly$\\alpha$ emission within color-selected Lyman Break Galaxy (LBG) populations}.  Tracking the redshift-dependent fraction in a well-defined population avoids consideration of absolute changes in the host galaxy number density such as has been the case in studies based on the luminosity function of narrow-band selected Lyman $\\alpha$ emitters (LAEs). Furthermore, evolution in dust obscuration can be independently tracked via correlations seen with the colors of the rest-frame UV continuum. Paper I presented a comprehensive survey of over 600 LBGs with deep spectra, mostly undertaken with the DEIMOS instrument on Keck, but including published samples from the Very Large Telescope (VLT, \\citet{Vanzella09} and references therein). In that paper we demonstrated the utility of the method and discussed the paucity of line emission in gravitationally-lensed $z>7$ candidates from the sample of \\citet{Richard08}.  In \\citet{Stark11}  (hereafter Paper II), through ultra-deep exposures with DEIMOS we significantly improved the line emission statistics at $z\\simeq 6$, providing a robust measure of the rest-frame EW distribution of Ly$\\alpha$ emission at the highest redshift when the Universe is believed to be fully ionized (\\citealt{Fan06}, c.f. \\citealt{Mesinger10}). This provides a sound basis for predicting the likelihood of emission at higher redshift and thereby enabling a test of whether there is absorption by neutral gas. Significantly, we found that over 50\\% of moderately-faint ($-20.25< \\rm{M}_{UV} <-18.75$) $z\\simeq 6$ LBGs exhibit strong emission with rest frame EWs $> 25$ \\AA. As this fraction increases over $4<z<6$, we argued on continuity grounds that we should expect a high success rate in recovering line emission from the newly-found WFC3/IR samples of $z>7$ LBGs {\\it unless} we encounter a more neutral IGM in the short time interval prior to $z\\simeq 7$.  Some evidence for this is seen in the recent studies of \\cite{Fontana10} and \\cite{Vanzella11}.  The present paper is concerned with an initial application of this test to the newly-available sample of WFC3/IR candidates with photometric redshifts in the redshift range $6.3<z<8.8$. Two important factors have motivated and shaped our program. Firstly, it is important to note that Ly$\\alpha$ emission is the only spectroscopic redshift indicator for galaxies beyond $z\\simeq 6$.  Since it is the absence of strong Ly$\\alpha$ emission that provides the basis for considering an increased neutral fraction, it is important to be sure that the targets are truly at the expected redshifts. Many early candidate LBGs believed to lie beyond $z\\simeq 6-7$ remained controversial because of their limited or marginal photometry.  The improved filter set and superior performance of WFC3/IR has given us confidence that the current list of $z\\simeq 7-8$ candidates is more robust than those based on earlier NICMOS data \\citep{Robertson10}.  Secondly, to match the lower redshift data, sampling from a similarly wide range of LBG luminosities, as we do here, will be advantageous.  As shown in Papers I and II, the fraction of line emission increases in intrinsically fainter sources and so by comparing fractions with respect to their LBG luminosities, we may gain additional evidence for the onset of the neutral era.  Throughout the paper, we adopt a $\\Lambda$-dominated, flat universe with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$ and $\\rm{H_{0}}=70\\,\\rm{h_{70}}~{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$. All magnitudes in this paper are quoted in the AB system \\citep{Oke83}. ",
        "conclusions": "Although we consider the most likely explanation for our observed decrease in the number of LBGs which show observable Ly$\\alpha$ emission to be an increase with redshift in the neutral fraction of the IGM, it is important to remember our assumptions. Foremost we have assumed that {\\it all} of our 26 targets have true redshifts beyond $z\\simeq 6.3$. Should there be low redshift interlopers or Galactic stars in our new sample, we will overestimate the decline in the Ly$\\alpha$ fraction. Secondly, we have assumed the DEIMOS spectra from Paper II constitute a representative sample for calculating the expected EW distribution for 6.3 $<z<$ 8.2. Although the uncertainties here are not as great, we plan further studies with DEIMOS to increase the statistical sample of 5.5 $<z<$ 6.3 LBGs.  Of course our observed decrease in the Ly$\\alpha$ fraction could also be attributed to an increased opacity arising from dust within the LBGs. However, given the blue UV continuum slopes observed for galaxies with $z>6.3$ \\citep{Bouwens10a,Dunlop11}, we consider this explanation unlikely.    Our diagnosis of a possible increase in the neutral hydrogen fraction beyond $z\\simeq$ 6.3 is supported by the earlier study of \\citet{Fontana10}. They found 1 marginal candidate out of 7 targets whereas we find 2 robust and 1 marginal cases out of our 19 targets spanning a larger luminosity and redshift range. Our conclusion is also supported by LAEs studies at $z=5.7$ and $6.5$ by \\cite{Ouchi10} and \\cite{Kashikawa11}. Compared to $z=5.7$, their LAE sample at  $z=6.5$ displays systematically lower EWs for Ly$\\alpha$.  They also derive little evolution in the rest UV luminosity function for LAEs, but a decrease in the Ly$\\alpha$ LF,  which could be explained by an increase in X$_{HI}$.  Our derived values of X$_{HI}$ are slightly higher than that of \\cite{Kashikawa11}, perhaps consistent with our survey probing to higher redshifts than their $z=6.5$ LAEs.  \\cite{Hayes11} have recently compiled results from numerous Ly$\\alpha$ and UV luminosity function studies to derive a volumetrically averaged Ly$\\alpha$ escape fraction, and find very similar results.  Their derived Ly$\\alpha$ escape fraction steadily increases with redshift below $z = 6$, then tentatively drops off at higher redshifts.  Very recently, \\cite{Ono11} report the convincing detection of Ly$\\alpha$ emission in a small fraction (3/11) of LBGs that, by virtue of their selection using Subaru imaging, are more luminous ($\\rm{M}_{UV}<-21$) than most of the objects considered here.  Such a complementary campaign targeting luminous LBGs selected from larger volumes will provide further insight into whether reionization is responsible for the declining fraction of line emission.  We note that our measured decrease in the fraction of LBGs with strong Ly$\\alpha$ potentially agrees with the result of \\cite{Cowie11}.  Although they argue against any evidence for reionization at $z = 6.5$, they find that  $\\sim24 \\%$ of galaxies at this redshift show strong Ly$\\alpha$ emission, comparable to the fraction we detect in this work, spread across a larger redshift range.  With the new generation of multi-object, near infrared spectrographs, such as MOSFIRE, set to come online soon, the prospects for this field are bright.  In addition to the significant multiplexing advantage, the increased sensitivity of these detectors will allow us to probe the lower luminosity ranges at $z\\geq 6.5$ to EW limits comparable to those in Paper II between sky lines.  Having such a statistical sample is key for allowing the quantification of any change in the hydrogen neutral fraction.  \\subsection*"
    },
    "1107/1107.3052_arXiv.txt": {
        "abstract": "{} {We present the results of flux density monitoring of PSR~B0329+54 at the frequency of 4.8~GHz using the 32-meter TCfA radiotelescope. The observations were conducted between 2002 and 2005. The main goal of the project was to find interstellar scintillation (ISS) parameters for the pulsar at the frequency at which it was never studied in detail. To achieve this the 20 observing sessions consisted of 3-minute integrations which on average lasted 24 hours. This gave us sufficient sensitivity to all types of flux density variations over a wide range of timescales. The character of the observations makes our project unique amongst other ISS oriented observing programs, at least at high frequency.} {Flux density time series obtained for each session were analysed using structure functions. For some of the individual sessions as well as for the general average structure function we were able to identify two distinctive timescales present, the timescales of diffractive and refractive scintillations. To the best of our knowledge, this is the first case when both scintillation timescales, $t_{\\rm DISS} = 42.7$ minutes and  $t_{\\rm RISS} = 305$ minutes, were observed simultaneously in a uniform data set and estimated using the same method.} {The obtained values of the ISS parameters combined with the data found in the literature allowed us to study the frequency dependence of these parameters over a wide range of observing frequencies, which is crucial for understanding the ISM turbulence. We found that the Kolmogorov spectrum is not best suited for describing the density fluctuations of the ISM, and a power-law spectrum with $\\beta =4$ seems to fit better with our results. We were also able to estimate the transition frequency (transition from strong to weak scintillation regimes) as $\\nu_c = 10.1$~GHz, much higher than was previously predicted. We were also able to estimate the strength of scattering parameter $u=2.67$ and the Fresnel scale as $r_F = 6.7 \\times 10^8$ meters.} {} ",
        "introduction": "Scintillations of the pulsar signal is a phenomenon that was identified shortly after the discovery of pulsars themselves (Hewish~et~al. \\cite{hewish}). From an observer point of view pulsars are point-like sources, and presumably are intrinsically stable sources of radiation. Their radiation, however, undergoes scattering process in the interstellar plasma on its way to the observer.  Depending on the strength of the scattering that the radiation of a given pulsar undergoes, the interstellar scintillations (ISS) will result in different phenomena observed for the pulsar (see Lorimer \\& Kramer, \\cite{lori}, and the references therein for a recent review). For a nearby pulsar, or when the radiation encounters only small amounts of the scattering material (and if the observing frequency is higher than ca. 1~GHz), the disturbances of the wavefront will be minimal, resulting in {\\it weak scintillations}.  In case of strong scattering (a distant pulsar and/or low observing frequency), the phase perturbations become large, which leads to a strong modulation of the pulsar signal. In this {\\it strong scintillation regime} one has to  take both diffraction and refraction effects into account. The {\\it diffractive interstellar scintillations} arise from small ($10^6 \\div 10^8$ meter) inhomogeneities in the interstellar medium, which introduce observed flux density variations with the timescale that can be estimated as $t_{\\rm DISS} \\sim f^{1.2} \\  d^{0.6}$ (assuming a Kolmogorov spectrum of the ISM turbulence, see Rickett \\cite{rick}; Romani~et~al. \\cite{romani}).  The {\\it refractive interstellar scintillations} (RISS; Sieber \\cite{sieb}; Rickett at al. \\cite{rick84}) arise from larger scale density fluctuations ($10^{10} \\div 10^{13}$ meters). As theories predict, the RISS timescale decreases with observing frequency, roughly as $t_{\\rm RISS} \\sim f^{-2.2} \\ d^{1.6}$.  For any given pulsar the values of the mentioned parameters will hence change with the observational frequency. Starting from very low frequencies, we will have short but increasing DISS timescales, and very long but decreasing RISS timescales, which will be accompanied by increasing RISS modulation (DISS modulation will be always close to unity). The rate at which these parameters change can be used to decide between various theories describing the ISM turbulence spectrum, because their frequency dependence power-law predictions are tied to the model used (see Romani~et~al. \\cite{romani}, Bhat~et~al. \\cite{bhat}).  As the observing frequency increases, both scintillation timescales should be closer, and the refractive modulation index should increase towards unity (which is best illustrated by Lorimer \\& Kramer, \\cite{lori}, on their Figure~4.5b). At a certain frequency, called the {\\it transition frequency} both timescales become equal, refractive scintillations reach maximum modulation and the character of the scintillations in general changes. This is the frequency at which the transition from the strong to weak scintillation regimes happens.  Most of the pulsar scintillation observations conducted so far were performed at the frequencies of 1.5~GHz and below, there is only a handful of papers describing ISS observations at higher frequencies (Malofeev~et~al. \\cite{malo}; Kramer~et~al. \\cite{kram}; Shishov et~al. \\cite{shis}). The attempts to estimate the spectrum of the ISM turbulence via the measurement of ISS parameters were limited to the frequency range from ca. $\\sim 100$~MHz to 1.5~GHz. This is both because pulsars are steep spectrum sources and are considerably weaker at higher frequencies, and the increasing DISS timescale which means one needs longer data sets to be able to perform measurements in a way similar to what is done at low frequencies. Both of these reasons result in a much longer observing time to be able to obtain reliable measurements. Because there are only few radiotelesopes capable of high-frequency observations, projects like this are very hard to implement.  The way to at least partially overcome these obstacles is to limit the observations to very strong pulsars, which in turn lowers the telescope-size requirements. One of the best candidates for these observations is definitely PSR~B0329+54, the strongest radio pulsar in the northern hemisphere. Owing to its brightness, it was observed by many various ISS projects in the past (see the {\\it Discussion} section in this paper, and Table~\\ref{tab3} for a full list of references). Amongst them was a project conducted by Malofeev~et~al.~(\\cite{malo}), which involved fairly short (ca. 130 minutes) observations at 4.75 and 10.55~GHz, based on which the authors concluded that for PSR~B0329+54 (with DM = 26.833) the transition frequency is ca. 3~GHz. This would mean that at higher frequencies this pulsar should be in the weak scintillation regime, and this statement was contradicting our personal experience with that source. The flux density of PSR~B0329+54 measured for the purposes of estimating pulsar spectrum (see Maron~et~al. \\cite{maro00} for an overview of that project) was showing very strong variations, and in individual measurements it ranged from a few to over a hundred milli-Janskys.  To solve this problem we decided to perform a long-time monitoring project of PSR~B0329+54 at the frequency of 4.8~GHz. We were kindly given a chance of extensive use of Toru\\'n Centre for Astronomy 32-meter radiotelescope (TCfA, located near Toru\\'n, Poland), which is equipped with a cooled 5~GHz receiver and the Penn State Pulsar Machine~II. We used it to perform numerous long-duration observing sessions between 2002 and 2005. This paper summarizes our results. ",
        "conclusions": "}  We presented the results of our three-year observing project of PSR~B0329+54 at the frequency of 4.8~GHz, using the 32-meter Toru\\'n Centre for Astronomy radiotelescope and the Penn State Pulsar Machine~II as a backend. The project was very unique, because it consisted of twenty separate observing sessions, each of them involving several hours (up to 5 days; more than 24 hours on average) of continuous observations of the flux density of the pulsar, with 3-minute time resolution. The total observing time ammounted to {\\it 35 thousand minutes}, or {\\it 24 days}, making it probably one of the longest observing programmes that involved long-term continuous observations of flux density and scintillation parameters that was performed using a single observing setup and uniformly analysed from raw data to the final results.  \\subsection{Scintillation parameters at 4.8 GHz}  The analysis of the data allowed us to conclude that flux density variations are governed by two distinctive timescales, which we identify as the diffractive and  refractive scintillation timescales. The fact that the observing frequency of 4.8 GHz is relatively close to the transition frequency for this pulsar makes those timescales differ only by a factor of $\\sim 7$. The character of our observations allowed us to detect them both at the same time, by means of the structure function analysis, using the same set of data - twice for individual sessions and additionally by constructing general average structure function. To our knowledge this is the first case of pulsar ISS observations where both scintillation timescales were measured from a single data set using exactly the same method.  The measured values of the diffractive timescale $t_{\\rm DISS}=42.7^{+4}_{-5}$~minutes and refractive timescale $t_{\\rm RISS}=305^{+40}_{-15}$~minutes allowed us to estimate the decorrelation bandwidth (unobtainable via direct measurements, because it exceeds the observing bandwidth of the backend used) as $B_{\\rm DISS}= 853$~MHz. Also, we were able to estimate the strength of scattering parameter, which can be expressed as square-root of the timescale ratio, as $u=2.67$.  We were also able to estimate the amount of modulation that both types of scintillation contribute to the flux variability, and did that by the means of the modulation indices. Especially important is the estimation of the modulation due to refractive scintillations, for which the modulation index is relatively strong: $m_{\\rm RISS}=0.56$. This agrees with the fact that during some of our observing sessions we noticed very drastic flux density variations, with the flux from single 3-minute integration reaching the values of 200~mJy, exceeding the average value by a factor of $\\sim 20$.  This shows how important it is to take ISS into account when measuring pulsar flux densities at high frequencies, especially for low-DM pulsars. For these objects the frequency of the switch from strong to weak scintillation regimes (i.e. the transition frequency) should be relatively low, and all theories predict the increase of refractive modulation index towards that frequency (for summary see Lorimer \\& Kramer \\cite{lori}; also Romani et~al. \\cite{romani}). In these cases, when the ISS timescales are of the order of tens to a few hundred minutes, typical flux density observations (which usually involve continuous integration of the length of 10 minutes to an hour) of a single measurement may be heavily affected by both types of scintillations. This in turn may cause problems when trying to construct the pulsar spectra, leading to artificial breaks or turn-ups observed in the spectrum - as was pointed out in a few cases by Maron~et~al.~(\\cite{maro00}). The only way around the problem is to repeat the measurement several times to average-out  the influence of ISS. Luckily, this should be less of a problem for pulsars with higher dispersion measures because the transition frequency for such objects should be well outside the usual observing frequencies, meaning that at least the refractive scintillations should not affect the measurement by much. Still, the diffractive scintillations would present a problem if their timescale would be similar to the integration lengths used, so one should approach every case carefully and individually.  Our measurements of the scintillation timescales allowed us also to estimate some of the parameters commonly used to describe the theory of interstellar scintillation. We were able to estimate the Fresnell scale $r_F$ as $6.7 \\times 10^8$~meter, which at the frequency of 4.8~GHz corresponds to a refractive scale of $s_r = 1.9 \\times 10^9$~m, and the diffractive timescale of $s_d=2.4\\times 10^8$~m.  We were also able to estimate the scintillation velocity of $V_{\\rm ISS} = 92.9$~km/s, and we hope that our value will help to solve the problem with the discrepancies that can be found in the literature. It agrees well with Wang~et~al.~(\\cite{wang}) as they obtained 97~km/s, on the other hand KS92 report $V_{\\rm ISS} = 68$~km/s, whereas Gupta~(\\cite{gupta95}) 170~km/s or 126~km/s, depending on the model. One has to remember that those discrepancies may arise because while the the formula is uniform (with the possible exception of the scaling factor $A_v$), the ways used to obtain the parameters were usually different.   \\subsection{General picture}  Adding our results to the estimations of the scintillation parameters at lower frequencies that can be found in the literature allowed us to study the dependence of these parameters on the observing frequency. That our observations were conducted at the frequency of 4.8~GHz proved to be very useful because it significantly widened the frequency range for which the values of ISS parameters are known. Previous attempts at this analysis were limited to the frequency of $\\sim 1.5$~GHz, adding our measurement expanded that range three-fold, which means an extension of ca. 50\\% when talking about the log-scale in frequency. This proved to be especially helpfull for the RISS parameters. At low frequencies the RISS timescale is very long and the modulation very weak, which brings a relatively large scatter to the low-frequency parameter estimates, making any attempts to fit a power-law type dependency questionable. The addition of our results, which we believe to be very reliable (because they come from a long-term monitoring programme), improves the general picture greatly (see Fig~\\ref{fig4}), making any fits to the data much more constrained.  We performed the fits of the power-laws governing the frequency dependence of the scintillation parameters, and they yielded us the respective spectral indices, which can be expressed as: $t_{\\rm DISS} \\sim f^{~1.01}$,  $t_{\\rm RISS} \\sim f^{~-1.76}$, $m_{\\rm  RISS} \\sim f^{~0.32}$ and $B_{\\rm DISS} \\sim f^{~3.68}$. These values significantly differ from those predicted by using Kolmogorov neutral gas theory to the ISM spatial electron density spectrum ($\\beta = 11/3$, see Table~\\ref{indx_table}). Indeed our measurements do not agree with any of the models describing the electron density spectrum that are commonly used, with the model using $\\beta = 4$ being the closest to the obtained spectral indices. This would indicate that the real spatial electron density spectrum is steeper than the Kolmogorov spectrum. Unfortunately, the character of our data did not allow for the estimation of the $\\beta$ index by the means used by a few other authors (see for example Wang~et~al.~\\cite{wang}) because they usually involve measurements performed in the dynamic spectra.  Using the above mentioned formulae we were able to estimate the transition frequency for PSR~B0329+54, i.e. the frequency at which the transition from strong to weak scintillation mode should appear, which is $\\nu_c = 10.1$~GHz. This clearly contradicts the earlier statement of Malofeev~et~al.~(\\cite{malo}), which estimated the transition frequency to be around 3~GHz. This value is also higher than the transition frequency one can infer from Lorimer \\& Kramer (\\cite{lori}) - namely from their Figure 4.3, which suggests that for a pulsar with $DM=27$~pc/cm$^3$ $\\nu_c$ should be close to 4~GHz.  It would be extremely interesting, but on the other hand very hard, to perform an observing project similar to ours at a frequency close to this transition frequency. First - the pulsar would be much weaker, thus requiring longer integration times. On the strong-scintillation-regime side of the transition frequency, the two ISS timescales would be extremely similar (100 minutes), which means that the distinction between them would be very hard to make, and would probably require very long sets of continuous observations. Because there are only a handful of telescopes that could perform these observations (i.e. have the required receiver equipment) this may be impossible to accomplish.  Nevertheless, any type of scintillation or flux density variation observations, performed at the frequency close to the transition frequency, may be very helpful to our understanding of the scintillation phenomenon in general, and the switch from strong- to weak-scintillation in particular. This is also true for other pulsars, especially those with lower dispersion measures, because the expected transition frequencies for them would be even lower than the 10~GHz for PSR~B0329+54."
    },
    "1107/1107.4720_arXiv.txt": {
        "abstract": "Discovering an Earth-like exoplanet in habitable zone is an important milestone for astronomers in search of extra-terrestrial life. While the radial velocity (RV) technique remains one the most powerful tools in detecting and characterizing exo-planetary systems, we calculate the uncertainties in precision RV measurements considering stellar spectral quality factors, RV calibration sources, stellar noise and telluric contamination in different observational bandpasses and for different spectral types. We predict the optimal observational bandpass for different spectral types using the RV technique under a variety of conditions. We compare the RV signal of an Earth-like planet in the habitable zone (HZ) to the near future state of the art RV precision and attempt to answer the question: How close are we to detecting Earth-like planet in the HZ using the RV technique? ",
        "introduction": "Surveys of exoplanets around stars of different spectral types using the RV technique have yielded fruitful discoveries including gas giants, Neptune-like planets and super-Earths. As of 2011 May, there are over 550 detected exoplanets and about 80\\% of them were discovered by the radial velocity (RV) technique using ground-based Doppler instruments\\footnotemark\\footnotetext{http://exoplanet.eu/; http://exoplanets.org/}. However, despite all these discoveries, an Earth-like planet in the HZ has not yet been detected.  We have not achieved the extremely high RV measurement precisions required in the detection of an Earth-like plaent in the HZ around stars of different masses. At lower end of stellar mass spectrum,~\\citet{Endl2006} measured the mean of RV RMS scatter for a sample of 90 M dwarfs, which is 8.3 $\\rm{m}\\cdot\\rm{s}^{-1}$. ~\\citet{Bean2010} demonstrated $\\sim$5 $\\rm{m}\\cdot\\rm{s}^{-1}$ RV precision for several M dwarfs over a period of 6 months. The intrinsic faintness and lack of an optimal RV calibration source currently limit the improvement of RV precision in the near infrared (NIR). In addition, imperfect stellar template spectra and modeling of telluric lines pose additional problems to extracting RV in the NIR where the flux of an M dwarf peaks~\\citep{Bean2010}. For solar type stars, i.e., spectral type of FGK, $\\sim$1 $\\rm{m}\\cdot\\rm{s}^{-1}$ RV precision has been routinely achieved~\\citep{Mayor2008, Howard2010}. However, stellar activity, stellar granulation and instrumental noise are major obstacles to overcome before detection of Earth-like planets~\\citep{Mayor2008}. At the high mass end, a precision of 3$-$6 $\\rm{m}\\cdot\\rm{s}^{-1}$ was demonstrated by ~\\citet{Johnson2007} and 6$-$8 $\\rm{m}\\cdot\\rm{s}^{-1}$ RV RMS scatter for stars with planets is found in the search of planets around evolved A type stars~\\citep{Bowler2010}. Fast stellar rotation broadens stellar absorption lines and limits the RV precision for main sequence massive stars~\\citep{Wright2005}. For evolved massive stars, similar problems exist as those facing solar type stars.  The fundamental photon-noise RV uncertainties have been discussed in several previous papers~\\citep{Butler1996,Bouchy2001}. However, only intrinsic properties of stellar spectra are discussed in their works while no detailed calculation of RV uncertainties introduced by the calibration sources. In a recent paper, ~\\citet{Reiners2010} considered the uncertainties in the NIR caused by RV calibration sources, i.e., a Th-Ar lamp and an Ammonia gas absorption cell. However, these two RV calibration sources can not be completely representative of the calibration sources used and proposed in current and planned Doppler planet survey in the NIR. For example, there are other emission lamps available in the NIR for RV calibration such as a U-Ne lamp as proposed by~\\citet{Mahadevan2010}.  In addition, other gas absorption cells besides the Ammonia cell have been proposed in the NIR~\\citep{Mahadevan2009, Valdivielso2010}.  Futhermore, in~\\citet{Reiners2010}, the calculation of RV calibration uncertainty of a gas absorption cell assumes a 50 nm band width in $K$ band, and then the uncertainty was applied to other NIR bands, which is purely hypothetical. Therefore, a more comprehensive and detailed study of RV calibration uncertainties is necessary at different observational bandpasses in the NIR in the search of planets around cool stars. On the other hand, in the visible, even though the current RV precision is not limited by the RV calibration source such as a Th-Ar lamp or an Iodine absorption cell, a better understanding of their performances under the photon-limited condition helps us discern a stage in which the RV calibration source becomes the bottle neck as RV precision keeps improving.~\\citet{Rodler2011} recently investigated RV precision achievable for M and L dwarfs, but did not quantitatively discussed the influence of RV calibration sources and stellar noise on precision Doppler measurement.  Stellar noise is a significant contributor to RV uncertainty budget, which falls into three categories: p-mode oscillation, spots and plagues, and granulations. P-mode oscillation usually produces an RV signature with a period of several minutes. The oscillation mode has been relatively well studied by previous work (e.g.,~\\citet{Carrier2003,Kjeldsen2005}). Exposure time of 10-15 min is proposed in order to smooth the RV signature induced by p-mode oscillation~\\citep{Dumusque2011}. Spots and plagues induced RV signal has been discussed by several papers (e.g.,~\\citet{Desort2007, Reiners2010, Lagrange2010, Meunier2010}).~\\citet{Meunier2010} concluded that the photometric contribution of plages and spots should not prevent detection of Earth-mass planets in the HZ given a very good temporal sampling and signal-to-noise ratio. Granulation is considered to be the major obstacle in detection of Earth planets in the HZ because it produces an RV signal with an amplitude of 8$\\sim$10 $\\rm{m\\cdot s}^{-1}$ based on observation on the Sun~\\citep{Meunier2010}. In addition, there is by far no good method of removing the RV noise from this phenomenon.~\\citet{Dumusque2011} provided a model of noise contribution in RV measurements based on precision RV observation on stars of different spectral type and at different evolution stages.  The telluric lines from the Earth's atmosphere are usually masked out in calculations of the photon-limited RV uncertainties in NIR~\\citep{Reiners2010, Rodler2011}. Although ~\\citet{Wang2011} proposed a method to quantitatively estimate the influence of atmosphere removal residual on precision Doppler measurement using the Dispersed Fixed Delay Interferometer (DFDI) method~\\citep{Ge2002,Erskine2003,VanEyken2010}, no attempt has ever been made for precision Doppler measurements using a high-resolution Echelle spectrograph. In practice, the telluric lines are not masked out, but instead modeled and removed. Therefore, a quantitative way of estimating the RV uncertainties produced by the residual of telluric line removal is necessary before we fully understand the performance of an RV instrument. In the visible band, the estimation of telluric line contamination is equally important as higher RV precision is required in the search of lower-mass planets around solar type stars.  We address two basic questions in this paper after considering a variety of factors including stellar spectrum quality, RV calibration precision, stellar noise and atmosphere contamination: 1, which observational bandpass is optimal to conduct precision Doppler measurements for stars of different spectral types; 2, is current RV precision adequate for detecting Earth-like planets in the HZ in the most optimistic scenario, in which the star is the least active and telluric lines are perfectly modeled and removed. The methods and findings of this study will provide insights to the design and optimization of a planned or ongoing precision Doppler planet survey. In addition, it also helps us to access at what stage we are in the search of Earth-like planets in the HZ.  The paper is organized as follows. In \\S \\ref{sec:method}, we introduce the methodologies of estimating the photon-limited RV uncertainties from difference sources. The findings are presented in \\S \\ref{sec:result}. Summaries and discussions of this study will be given in \\S \\ref{sec:Conclusion}. ",
        "conclusions": "\\label{sec:Conclusion}  We provide a method of practically estimating the photon-limited RV precision based on the spectral quality factor, stellar rotation, RV calibration uncertainty, stellar noise and telluric line contamination. The methodology described and the results presented in this paper can be used for design and optimization of planned and ongoing precision Doppler planet surveys. For pure consideration of stellar spectral quality without artificial rotationally broadening the absorption line profile, the optimal band for RV planet search is $B$ band for a wide range of spectral types from K to A, while it is $R$ or $K$ band for mid-late type M dwarfs. Nevertheless, the above conclusion remains unchanged after considering typical stellar rotation of each spectral type. However, F and A stars become unsuitable for precision RV measurements because of typically fast stellar rotation. We confirm the finding in~\\citet{Reiners2010} that the NIR Doppler measurements gain advantage for mid-late M dwarfs. However, instead of finding $Y$ band as the optimal band considering stellar spectrum quality and telluric masking, we find that both $Y$ and $H$ bands give the highest RV precision among $V$, $Y$, $J$ and $H$ bands. In a comparison to~\\citet{Rodler2011}, we find $K$ band is the optimal band for precision Doppler measurement in a telluric-free case and $H$ band is optimal in a telluric-masking case, while they found $Y$ band gives the highest RV precision in both cases. Fundamental photon-limited RV precision for evolved stars has been discussed by~\\citet{Jiang2011}, which is valuable for ongoing RV planet search around retired stars discussed in~\\citet{Johnson2007}.  We also consider the uncertainties brought by current available RV calibration sources at different spectral resolutions (Fig. \\ref{fig:Un_Cal_Res}). Sub $\\rm{m}\\cdot\\rm{s}^{-1}$ calibration precision can be reached for each observational bandpass. Note that the Q factors may change as gas pressure, length of light path and temperature changes. The precision also depends on the methods used in the RV calibration. We categorized the calibration methods into several cases: Superimposing, in which the calibration spectrum is imprinted onto a stellar spectrum; Non-Common Path and Bracketing, in which the calibration is conducted either spatially or temporally. The former method depends on stellar flux while the latter one can only be applicable for very stable instruments. There are other calibration sources we have not included into the discussions in this study, for example, laser combs~\\citep{Steinmetz2008,Li2008}, the Fabry-Perot calibrator~\\citep{Wildi2010} and the Monolithic Michelson Interferometer~\\citep{Wan2010}. Once they become more economically affordable or more technically ready, the RV precision will be greatly improved in the future.  For the first time we have quantitatively estimated the uncertainty caused by the residual of telluric contamination removal for high resolution echelle spectroscopy method. Depending on the telluric absorption, different observational bandpasses are affected differently. $B$ band is the least sensitive to telluric contamination because there are barely any telluric absorption features in $B$ band. However, the NIR bands suffer the most in precision RV measurements because the stellar absorption lines and telluric lines are mixed together severely in this spectral region. Only when $\\alpha\\leq$0.01, i.e., more than 99\\% of strength of telluric lines is removed, the advantage of NIR observation of mid-late type M dwarfs begins to show, which is a factor of 3 improvement. This quantitative method in estimating the RV uncertainty induced by telluric contamination can be easily adapted to other problems, for example, estimating the moon light contamination.  Besides telluric line removal, telluric line masking has also been discussed in several of previous studies~\\citep{Reiners2010, Rodler2011, Wang2011}. In ~\\citet{Reiners2010}, telluric absorption with depth more than 2\\% and 30 $\\rm{km}\\cdot\\rm{s}^{-1}$ in the vicinity is blocked out when measuring RV. Based on this blocking criterium, the photon-limited RV uncertainty, $\\delta v_{rms, S}$(refer to Equation (\\ref{eq:simple_example_tulleric_2})) , for an M9V star at $\\rm{R}$=100,000 is 3.9, 2.2, 3.9, 2.2 $\\rm{m}\\cdot\\rm{s}^{-1}$ in $V$, $Y$, $J$ and $H$ band respectively (see Table \\ref{tab:Comp_Reiners}). In comprison, $\\delta v_{rms, N}$ is 71.3, 5.8, 6.5, 3.7 $\\rm{m}\\cdot\\rm{s}^{-1}$ in $V$, $Y$, $J$ and $H$ band respectively. Except for $V$ band, $\\delta v_{rms, S}$ and $\\delta v_{rms, N}$ are at the same order of magnitude, and the uncertainty caused by telluric absorption cannot be neglected even though that the spectral region with any telluric absorption of more than 2\\% is blocked. If more strict criterium of telluric line masking is applied, fewer photons are considered in measuring the RV, which effectively increases the photon-limited RV uncertainty. In order to reach photon-limited RV precision predicted by pure consideration of spectral Q factor, telluric removal should be applied in which telluric contamination is measured or modeled and then removed from measured stellar spectrum.  RV uncertainty due to stellar granulation is taken into consideration in this paper. High frequency ($\\sim$min) stellar noise such as p-mode oscillations usually have a RV amplitude of 0.1 to 4.0 $\\rm{m}\\cdot\\rm{s}^{-1}$ ~\\citep{Schrijver2000} and they can be averaged out within typical 10$-$15 exposure time. RV uncertainties due to low frequency ($10-100$ day) stellar noise such as stellar spots have been discussed in recent papers, for example, ~\\citet{Desort2007} and ~\\citet{Reiners2010}. The amplitudes of spot-induced RV range from one to several hundred $\\rm{m}\\cdot\\rm{s}^{-1}$. Since stellar spot induced RV uncertainties are periodic and therefore can be modeled and removed, however, the amplitude of residual is unknown at this stage.  We compare the RV precision based on stellar spectral quality and the signal of an Earth-like planet in the HZ of a star with a certain spectral type. We find that it is likely to detect a habitable Earth-like planet around G,K and M stars while it is too demanding to detect one around F and A stars. $B$ band is the optimal band for G and K stars and $K$ band for M dwarfs. After considering practical issues such as telluric contamination, we find that, except for $B$ band, every observational bandpass is affected by telluric contamination to some extent. The major RV measurement uncertainty comes from telluric contamination, which overwhelms the RV signal of an habitable Earth-like planet around G and K stars. Surprisingly, telluric contamination becomes an issue in $V$ band even there is only 2.4\\% of spectral region affected by telluric lines. After telluric lines are removed at a very high level, i.e., $\\alpha\\leq$0.001, the error from RV calibration becomes the major contributor of Doppler measurement uncertainty. After stellar noise (granulation only) is taken into consideration, which is dominant contributor to RV uncertainty, M dwarfs become the only type of star that is suitable for the search for Earth-like planets in the HZ.  The RV precision in the discussion of habitable Earth detectability considers four factors: stellar spectral quality, RV calibration uncertainty, stellar noise and telluric contamination. However, the discussion of stellar noise should be treated with great caution for several reasons: 1, stellar noise is not very well understood and characterized at this stage; 2, it is different from case to case and therefore it is difficult to draw a general conclusion; 3, a habitable planet search is different from a planet survey, the targets are chosen in favor of detection at the best case scenario, for example, high stellar flux, slow stellar rotation, low stellar activity and low stellar noise and so on. Therefore, the stellar noise assumed in this study is in the best case scenario according to current theory and observation. In addition, we assume the highest signal within linear range (30,000 ADU) for current typical CCD (16-bit dynamic range) in single exposure in the discussion, note that the S/N can also be improved by multiple independent measurements. The S/Ns required for Earth-like planet detections are provided in Table \\ref{tab:SNR_SP} based on stellar spectral quality. Please note that the HZ changes over time as the luminosity of the host star changes. It also depends on properties of a planet such as atmosphere composition, albedo and orbit. The purpose of discussion in this paper is to provide a basic idea of the comparison of current best obtainable RV precision to a typical RV signal of an habitable Earth-like planet.  We  acknowledge the support from NSF with grant NSF AST-0705139, NASA with grant NNX07AP14G (Origins), UCF-UF SRI program, DoD ARO Cooperative Agreement W911NF-09-2-0017, Dharma Endowment Foundation and the University of Florida. We thank Stephen Redman for kindly providing us the lines list of a U-Ne emission lamp. We thank Suvrath Mahadevan, Cullen Blake, Peng Jiang and Mark Keremedjiev for valuable discussion and comments."
    },
    "1107/1107.3877_arXiv.txt": {
        "abstract": "The NGC 1333 IRAS 4A protobinary was observed in the 1.3 cm and 6.9 mm continuum and the ammonia and SiO lines, with an angular resolution of about 0.4 arcseconds. The continuum maps show the circumstellar structures of the two protostars, A1 and A2. The A1 system is brighter and more massive than the A2 system. The ratio of mass, including dense gas and protostar, is about 6. The properties of the circumstellar disks and outflows suggest that A1 may be younger than A2. The deflected part of the northeastern jet of A2 is bright in the SiO line, and the distance between the brightest peak and deflection point suggests that the enhancement of SiO takes about 100 yr after the collision with a dense core. The ammonia maps show a small structure that seems to be a part of the obstructing core. The outflow properties were studied by comparing interferometric maps of SiO, ammonia, formaldehyde, and HCN lines. Their overall structures agree well, suggesting that these species are excited by the same mechanism. However, the intensity distributions show that SiO is chemically unique. SiO may be directly linked to the primary jet while the other species may be tracing the entrained ambient gas. ",
        "introduction": "NGC 1333 IRAS 4A2 is one of the best-studied low-mass protostars. It has a circumstellar disk viewed nearly edge-on and drives a well-collimated and rotating outflow/jet (Choi 2005; Choi et al. 2007, 2011). Choi et al. (2010) found that the NH$_3$ (2, 2) and (3, 3) lines seem to selectively trace the accretion disk while the millimeter/centimeter continuum seems to trace both disk and envelope. The rotation kinematics of the disk suggests that the mass of the central protostar is $\\sim$0.08 $M_\\odot$ and that the collapse age is $\\sim$50,000 yr (Choi et al. 2010). The bipolar jet flows in the northeast-southwest direction (Blake et al. 1995; Hodapp \\& Ladd 1995; Lefloch et al. 1998). The jet may be launched from a small region on the disk, or outflow foot-ring, with a radius of $\\sim$2 AU (Choi et al. 2011). The northeastern jet shows a sharp bend of flow direction that may be caused by a collision with a dense core in the ambient cloud (Choi 2005; Baek et al. 2009). IRAS 4A2 also exhibits other star formation activities such as H$_2$O maser emission (Park \\& Choi 2007; Marvel et al. 2008).  IRAS 4A2 belongs to a binary system (IRAS 4A) in the Perseus star-forming region at a distance of 235 pc from the Sun (Lay et al. 1995; Looney et al. 2000; Hirota et al. 2008). The spectral energy distribution of IRAS 4A suggests that the constituent binary members are Class 0 protostars (Sandell et al. 1991; Enoch et al. 2009). The luminosity of IRAS 4A suggests that the protostars are growing at a rate typical of Sun-like stars through the accretion of gas from the molecular envelope (Jennings et al. 1987; Choi et al. 2010).  IRAS 4A1 is brighter in the millimeter/centimeter continuum than its sibling protostar A2 (Looney et al. 2000; Reipurth et al. 2002; J{\\o}rgensen et al. 2007), but it is weaker in the NH$_3$ emission (Choi et al. 2007). IRAS 4A1 seems to drive an outflow to the south, but its counter flow, if any, has not been detected clearly (Choi 2005). While A1 must be in an evolutionary stage similar to that of A2, the details are less clear.  In this paper, we present the results of our observations of the NGC 1333 IRAS 4A region in the 1.3 cm and 6.9 mm continuum, the NH$_3$ lines, and the SiO line with an angular resolution higher than those of the previous studies (Choi 2005; Choi et al. 2007). We describe our observations in Section 2. In Section 3, we report the results and discuss the star-forming activities in the IRAS 4A region. A summary is given in Section 4. ",
        "conclusions": ""
    },
    "1107/1107.5456_arXiv.txt": {
        "abstract": "{ We analyze the long-term evolution of the fluxes of six active galactic nuclei (AGN) -- 0923+392, 3C~111, 3C~273, 3C~345, 3C~454.3, and 3C~84 -- in the frequency range 80--267~GHz using archival calibration data of the IRAM Plateau de Bure Interferometer. Our dataset spans a long timeline of $\\approx$14 years with 974 -- 3027 flux measurements per source. We find strong (factors $\\approx$2--8) flux variability on timescales of years for all sources. The flux density distributions of five out of six sources show clear signatures of bi- or even multimodality. Our sources show mostly steep ($\\alpha\\approx0.5-1$), variable spectral indices that indicate outflow dominated emission; the variability is most probably due to optical depth variations. The power spectra globally correspond to red-noise spectra with five sources being located between the cases of white and flicker noise and one source (3C~111) being closer to the case of random walk noise. For three sources the low-frequency ends of their power spectra appear to be upscaled in spectral power by factors $\\approx$2--3 with respect to the overall powerlaws. In two sources, 3C~454.3 and 3C~84, the 1.3-mm emission preceeds the 3-mm emission by $\\approx$55 and $\\approx$300 days, respectively, probably due to (combinations of) optical depth and emission region geometry effects. We conclude that the source emission cannot be described by uniform stochastic emission processes; instead, a distinction of ``quiescent'' and (maybe multiple) ``flare'' states of the source emission appears to be necessary. } ",
        "introduction": "Active galactic nuclei (AGN) have been studied extensively in the wavelength range from cm-radio to $\\gamma$-rays in the last decades (see, e.g., Kembhavi \\& Narlikar \\cite{kembhavi1999}, or Krolik \\cite{krolik1999}, and references therein for a review). There is overwhelming observational evidence that their emission originates from accretion onto supermassive black holes (SMBH) with masses $M_{\\bullet}\\simeq10^{6...9}M_{\\odot}$ (e.g., Ferrarese \\& Ford \\cite{ferrarese2005}, and references therein).  The viewing-angle unification standard model of AGN (Lawrence \\cite{lawrence1987}; see also Urry \\& Padovani \\cite{urry1995}) has long been established, although recent studies emphasize the impact of galaxy mergers on AGN fueling and luminosity (e.g. Hopkins et al. \\cite{hopkins2006}). However, the details of AGN activity are still not well understood. One example is the nuclear emission; proposals for its origin range from shocks in continuous (e.g., Marscher \\& Gear \\cite{marscher1985}) or discontinuous (e.g., Spada et al. \\cite{spada2001}) jets to orbiting plasma ``hotspots'' (e.g., Abramowicz et al. \\cite{abramowicz1991}) or plasma density waves (e.g., Kato \\cite{kato2000}) in accretion disks.  One possibility to constrain the AGN physics is offered by the analysis of characteristic timescales and the noise properties of the emission. Several studies aimed at identifying characteristic timescales, including possible quasi-periodic oscillations, from tens of minutes (using X-ray observations; e.g. Benlloch et al. \\cite{benlloch2001}, and using mm-radio observations; e.g. Sch\\\"odel et al. \\cite{schoedel2007}) to tens of years (based on cm- and mm-radio monitoring data; Hovatta et al. \\cite{hovatta2007,hovatta2008}). Other works focused on the power spectral statistics of the emission which was found to globally follow red-noise type powerlaws (e.g. Lawrence \\& Papadakis \\cite{lawrence1993}; Uttley et al. \\cite{uttley2002}).  The studies by Hovatta et al. (\\cite{hovatta2007,hovatta2008}) used radio data in the frequency range 22--87~GHz spanning timelines of $\\approx$25 years. A key result of their study was that AGN show large-amplitude activity on timescales $\\gtrsim$6 years. Therefore, analysing the long-term behaviour of AGN requires monitoring on timescales of $\\gtrsim$10 years. Building up on this conclusion, we have analyzed archival flux monitoring data from the IRAM Plateau de Bure Interferometer (PdBI) of six AGN -- 0923+392, 3C~111, 3C~273, 3C~345, 3C~454.3, and 3C~84 -- in the (observatory frame) frequency range 80--267~GHz. Given the range of redshifts $z\\approx0.02-0.86$ of our sample, this corresponds to a range of rest-frame frequencies of $\\approx$82--497~GHz. Our dataset comprises several thousand measurements for each source, covering timelines of $\\approx$14 years. We note that our study complements earlier AGN mm/radio surveys carried out with the IRAM Pico Veleta observatory (Steppe et al. \\cite{steppe1988,steppe1992,steppe1993}; Reuter et al. \\cite{reuter1997}). ",
        "conclusions": "We have analyzed the long-term activity of six luminous active galactic nuclei using mm/radio lightcurves spanning $\\approx$14 years in time. From our analysis we draw the following principal conclusions:  \\begin{enumerate}  \\item  All sources show substantial, non-periodic variability by factors $\\approx$2--8 on timescales of years. This is in good agreement with the results reported by Hovatta et al. (\\cite{hovatta2007,hovatta2008}) and shows that source monitoring times on scales of decades are necessary in order to obtain a realistic overview on the source activity. We also note that at least for the special case of Sagittarius~A* both observational (e.g. Koyama et al. \\cite{koyama2008}) and theoretical (e.g. Cuadra et al. \\cite{cuadra2008}) evidence suggests activity timescales of several hundred years. An even more extreme example is provided by recent observations of ``Hanny's Voorwerp'' which suggest flux variations of the galactic nucleus of IC~2497 by factors $>$100 on timescales $<$70,000 years (e.g., Schawinski et al. \\cite{schawinski2010}). Therefore one should be aware that AGN monitoring studies, like our work, might observe some of their targets at special times even when covering timelines of decades.  \\item  The flux density distributions of five out of six sources show clear signatures of bi- or even multimodality. This indicates that the emission processes cannot be treated as uniform stochastic processes. Instead, we have to assume a segregation of ``quiescent'' and (maybe multiple) ``flare'' source states that correspond to different physical regimes. Each state would be characterized by an individual (probably red-noise type) flux distribution.  \\item  Our sources show mostly steep ($\\alpha\\approx0.5-1$) spectral indices that indicate outflow (halos, jets) dominated synchrotron emission. The spectral indices show substantial variability that is not correlated with the source fluxes. From the combined information we conclude that optical depth variations are the dominant mechanism for the spectral variability.  \\item  All sources show power spectra that are globally consistent with red noise (i.e. $\\beta>0$). We find that the intrinsic (i.e. measurement noise corrected) power-spectral indices are flatter than 1 for five of our sources, placing their emission between the cases of flicker noise ($\\beta=1$) and white noise ($\\beta=0$). The only exception is 3C~111 whose power spectrum is consistent with being located between flicker noise and random walk noise ($\\beta=2$). For three sources we find that the low-frequency ends of their power spectra appear to be upscaled in spectral power by factors $\\approx$2--3 with respect to the overall powerlaws, presenting another indication towards a segregation between ``quiescent'' and ``flaring'' source states.  \\item  For four of our targets -- 0923+392, 3C~111, 3C~273, and 3C~345 -- we find that the 1.3-mm and 3-mm lightcurves are simultaneous within $3\\sigma$ confidence levels. This is in agreement with shock-models like the ``generalized'' model by Valtaoja et al. (\\cite{valtaoja1992}); given the measurement accuracies of tens of days however these constraints are fairly loose. For 3C~454.3 (in its ``flare state(s)'' after 2005.2) and 3C~84 we find spectral time delays of $\\approx$55 days and $\\approx300$ days, respectively, with the 1.3-mm emission leading. The behaviour of 3C~454.3 can be traced back to variations of the optical depth with time and frequency, in agreement with previous work. In case of 3C~84, we suspect a combination of optical depth and emission region geometry effects as cause for the time lag.  \\item  We do not see a correlation between the observed source properties -- variability, spectral index, power spectrum -- with physical source parameters -- redshift, black hole mass, kilo-parsec morphology -- as given in Table~\\ref{tab_journal}.  \\end{enumerate}  The discovery of distinct source emission states challenges the common assumption that AGN emission can be described by uniform stochastic processes. A promising road for further studies is indicated by the cases of microquasars, Sagittarius~A*, and IC~2497 where similar segregations of emission states have been seen. Long-term, multi-wavelength monitoring of AGN appears to be the most important tool towards a deeper understanding of this phenomenon."
    },
    "1107/1107.0335_arXiv.txt": {
        "abstract": "{The {\\it Kepler} spacecraft is providing time series of photometric data with micromagnitude precision for hundreds of A-F type stars.}{We present a first general characterization of the pulsational behaviour of A-F type stars as observed in the {\\it Kepler} light curves of a sample of 750 candidate A-F type stars, and  observationally investigate the relation between $\\gamma$\\,Doradus (\\gdor), $\\delta$\\,Scuti (\\dsct), and hybrid stars.}{We compile a database of physical parameters for the sample stars from the literature and new ground-based observations. We analyse the {\\it Kepler} light curve of each star and extract the pulsational frequencies using different frequency analysis methods. We construct two new observables, {\\it 'energy'} and {\\it 'efficiency'}, related to the driving energy of the pulsation mode and the convective efficiency of the outer convective zone, respectively.}{We propose three main groups to describe the observed variety in pulsating A-F type stars:  \\gdor, \\dsct, and hybrid stars.  We assign 63\\% of our sample to one of the three groups, and identify the remaining part as rotationally modulated/active stars, binaries, stars of different spectral type, or stars that show no clear periodic variability. 23\\% of the stars (171 stars) are hybrid stars, which is a much higher fraction than what has been observed before. We characterize for the first time a large number of A-F type stars (475 stars) in terms of number of detected frequencies, frequency range, and typical pulsation amplitudes. The majority of hybrid stars show frequencies with all kinds of periodicities within the \\gdor\\,and \\dsct\\,range, also between 5 and 10 d$^{-1}$, which is a challenge for the current models. We find indications for the existence of \\dsct\\,and \\gdor\\,stars beyond the edges of the current observational instability strips. The hybrid stars occupy the entire region within the \\dsct\\,and \\gdor\\,instability strips and beyond.  Non-variable stars seem to exist within the instability strips. The location of \\gdor\\,and \\dsct\\,classes in the (\\Teff, \\logg)-diagram has been extended. We investigate two newly constructed variables, {\\it 'efficiency'} and {\\it 'energy'}, as a means to explore the relation between \\gdor\\,and \\dsct\\,stars.}{Our results suggest a revision of the current observational instability strips of \\dsct\\,and \\gdor\\,stars and imply an investigation of pulsation mechanisms to supplement the $\\kappa$\\,mechanism and convective blocking effect to drive hybrid pulsations. Accurate physical parameters for all stars are needed to confirm these findings.} ",
        "introduction": "With the advent of the asteroseismic space missions MOST (Walker et al. 2003), CoRoT  (Baglin et al. 2006), and {\\it Kepler} (Borucki et al. 2010), a new window is opening towards the understanding of the seismic behaviour of A- and F-type pulsators. The main advantages of these space missions are (1) the long-term continuous monitoring of thousands of stars, which enables both the determination of long-period oscillations and the resolving of beat frequencies; and (2) the photometric precision at the level of milli- to micro-magnitudes, which will provide a more complete frequency spectrum and also allow the detection of low-amplitude variations that are unobservable from the ground and providing a more complete frequency spectrum.  The availability of these long-term, very precise light curves makes possible the first comprehensive analysis of the variability of a sample of several hundred candidate A-F type stars that is presented here.   The region of variable A- and F-type, including main sequence (MS), pre-MS, and post-MS stars,  with masses between 1.2 and 2.5 M$_{\\odot}$ hosts the $\\gamma$\\,Doradus (\\gdor) and $\\delta$\\,Scuti (\\dsct) pulsators. The \\gdor\\,stars were recognized as a new class of pulsating stars less than 20 years ago (Balona, Krisciunas \\& Cousins 1994).  Our current understanding is that they pulsate in high-order gravity (g) modes (Kaye et al. 1999a), excited by a flux modulation mechanism induced by the upper convective layer (Guzik et al. 2000; Dupret et al. 2004; Grigahc\\`ene 2005). Typical \\gdor\\,periods are between 8~h and 3~d. From the ground, about 70 {\\it bona fide} and 88 candidate \\gdor\\,pulsators have been detected (Balona et al. 1994; Handler 1999; Henry, Fekel \\& Henry 2005; De Cat et al. 2006; Henry et al. 2011; among other papers).  The \\dsct\\,variables, on the other hand, have been known for decades. They show low-order g and pressure (p) modes with periods between 15\\,min and 5\\,h that are self-excited through the  $\\kappa$-mechanism (see reviews by Breger 2000; Handler 2009a). Several hundreds of \\dsct\\,stars have been observed from the ground (e.g. catalogue by Rodr\\'{\\i}guez \\& Breger 2001).  Because the instability strips of both classes overlap, the existence of hybrid stars, i.e. stars showing pulsations excited by different excitation mechanisms, is expected, and a few candidate hybrid stars have indeed been detected from the ground (Henry \\& Fekel 2005;  Uytterhoeven et al. 2008; Handler 2009b).   The main open question in seismic studies of A- and F-type stars concerns the excitation and mode selection mechanism of p and g modes. The only way to understand and find out systematics in the mode-selection mechanism is a determination of pulsation frequencies and pulsation mode parameters for a large number of individual class members for each of the pulsation classes, and a comparison of the properties of the different case-studies. So far, a systematic study of a sufficiently substantial sample was hampered by two factors. First, the number of detected well-defined pulsation modes is too small  to construct unique seismic models, which is caused by ground-based observational constraints, such as bad time-sampling and a high noise-level. Second, only a small number of well-studied cases exist, because a proper seismic study requires a long-term project, involving ground-based multi-site campaigns spanning several seasons, or a dedicated space mission.   First demonstrations of the strength and innovative character of space data with respect to seismic studies of A-F type stars are the detection of two hybrid \\gdor-\\dsct\\,stars by the MOST satellite (HD\\,114839, King et al. 2006; BD+18-4914, Rowe et al. 2006), and the detection of an impressive number of frequencies at low amplitudes, including high-degree modes as confirmed by ground-based spectroscopy, in the precise space CoRoT photometry of the \\dsct\\,stars HD\\,50844 (Poretti et al. 2009) and HD\\,174936 (Garc\\'{\\i}a Hern\\'andez et al. 2009), and the \\gdor\\,star HD\\,49434 (Chapellier et al. 2011). The first indications that hybrid behaviour might be common in A-F type stars were found from a pilot study of a larger sample of {\\it Kepler} and CoRoT stars (Grigahc\\`ene et al. 2010; Hareter et al. 2010). Recently, Balona et al. (2011a) announced the detection of \\dsct\\,and \\gdor\\,type pulsations in the {\\it Kepler} light curves of Ap stars. Hence,  a breakthrough is expected in a currently poorly-understood field of seismic studies of A-F type pulsators through a systematic and careful investigation of the pulsational behaviour in a large sample of stars.  The goals of the current paper are (1) to present a first general characterization of the pulsational behaviour of main-sequence A-F type stars as observed in the {\\it Kepler} light curves of a large sample; and (2) to observationally investigate the relation between \\gdor\\,and \\dsct\\,stars and the role of hybrids. In forthcoming papers, detailed seismic studies and modelling of selected stars will be presented. ",
        "conclusions": "We analysed the {\\it Kepler} light curves based on survey phase data with time spans between 9\\,d and 322\\,d available through KASOC and associated frequency spectra of 750 candidate A-F type stars in search for \\dsct, \\gdor, and hybrid pulsators. The main results are:  \\begin{itemize} \\item The {\\it Kepler} light curves of the sample of 750 candidate A-F type stars show a variety in variability behaviour. \\item Observationally, we propose three main groups to describe the observed variety:  \\gdor, \\dsct, and hybrid stars. The latter group includes both \\dsct-dominated and \\gdor-dominated hybrid stars. About 63\\% of the sample are unambiguously assigned to one of the three groups. \\item About 23\\% of the sample are hybrid candidates (171 stars, or 36\\% of the stars assigned to the three groups). This is in strong contrast with the number of hybrid candidates so far observed from the ground, but compatible with the first {\\it Kepler} study of \\gdor\\,and \\dsct\\,variables by Grigahc\\`ene et al. (2010).  The far superior precision of the {\\it Kepler} space data opens a new window in detecting low-amplitude variations. {\\it Kepler} will be ideal to study hybrid behaviour in different types of stars, such as roAp stars (Balona et al. 2011a), sdB stars (\\O{}stensen et al. 2010), and B stars (Balona et al. 2011b). \\item We presented a characterization of the stars in terms of number of detected frequencies, frequency range, and typical pulsation amplitudes, which provides valuable feedback for models and instability studies. This is the first time that this kind of information is available for a substantial sample of stars. Up to 500 non-combination frequencies are detected in the {\\it Kepler} time series of a single star. The highest pulsation amplitude measured is 58\\,000\\,ppm. The shortest detected \\dsct\\,periods are about 18\\,min. We find that hybrid stars show all kinds of periodicities within the \\gdor\\,and \\dsct\\,range. In particular, the majority of hybrid stars shows frequencies between 5 and  10\\,d$^{-1}$. From a theoretical point of view, this result presents a number of challenges, because the currently accepted over-stability mechanisms cannot explain the presence of pulsational modes in the wide frequency ranges observed with {\\it Kepler}. It needs to be investigated if and to what extent the presence of stochastic modes, high-degree, and/or rotationally split modes with high amplitudes, granulation and effects of convection can explain part of the unexpected observed modes. \\item The location of \\gdor\\,and \\dsct\\,classes in the (\\Teff, \\logg)-diagram has been extended (Fig.\\,\\ref{logTefflogg}).  We find indications that {\\it Kepler} \\dsct\\,stars exist beyond the red edge of the observational instability strip, while {\\it Kepler} \\gdor\\,pulsations seem to appear in both hotter and cooler stars than observed so far. The {\\it Kepler} hybrid stars occupy the entire region between the blue edge of the \\dsct\\,instability strip and the red edge of the \\gdor\\,instability strip and beyond.  These results, if confirmed by verification of the temperature and \\logg\\,values in a more comprehensive sample, imply that the observational instability strips need to be extended to accommodate the {\\it Kepler} \\dsct\\,and \\gdor\\,stars. From a theoretical point of view, the overall presence of hybrid stars implies an investigation of other pulsation mechanisms to supplement the $\\kappa$\\,mechanism and convective blocking effect to drive hybrid pulsations. \\item Two new 'observables' that reflect the different internal physics of \\dsct\\,and \\gdor\\,pulsators are introduced to investigate the relation between the two types of pulsations (Fig.\\,\\ref{diagram}): (1)  {\\it efficiency},  related to the convective efficiency of the outer convective zone, and a function of \\Teff\\,and \\logg;  and (2) {\\it energy}, the driving energy of a mode, and a function of the highest observed frequency amplitude and the associated frequency. Both observables are empirical and are constructed using only available measured variables. The impact and physical significance of the group separation in the ($\\log$({\\it efficiency}), $\\log$({\\it energy}))-diagram needs to be  investigated in more detail. The two new observables are a promising starting point for further investigations of the relation between \\dsct, \\gdor\\,and hybrid stars. \\item Our study indicates that Kp\\,$= 14$\\,mag is a cut-off magnitude for detection of variations with amplitudes below 20\\,ppm in A-F type stars with {\\it Kepler}. \\item Sixteen percent of the sample stars show no clear variability within the expected range of frequencies for \\dsct\\,and \\gdor\\,stars. Faint and cool stars predominate this sample. Among the stars, we identified 75 candidate solar-like stars. No correlation between  non-variability and the length of the available dataset or the available cadence mode is found. We find indications for the presence of constant stars inside the instability strips of A-F type pulsators. \\item The remaining 21\\% of sample stars are identified as a Cepheid, B-type stars, red giant stars, stars that show stellar activity, or binaries. At least 12\\% of the sample are identified as a binary or multiple system, based on investigation of the {\\it Kepler} light curve or on input from the literature. Many long-period binaries are expected to be among the remaining stars of the sample.  3.5\\% of the sample stars shows eclipses. Several of the EBs have variable components, including  \\dsct, \\gdor, and hybrid stars. \\end{itemize}   Clearly, space missions are changing the landscape of  \\gdor\\,and \\dsct\\,pulsators. We aimed at a global analysis of the sample stars. A careful seismic analysis of individual stars is needed to confirm their classification,  clarify the observed variety in pulsational behaviour, fully characterize the properties of the \\dsct, \\gdor, and hybrid groups,  understand their relationship, clarify the driving mechanism(s) for each group,  and elaborate on the variables {\\it energy} and {\\it efficiency}. The observational results with {\\it Kepler} presented here open up several new questions and theoretical challenges for the current models related to pulsational instability, thermodynamics, and stellar structure. We mention here some topics for further investigation.  To be able to place the stars confidently in the (\\Teff, \\logg)-diagram,  estimate the projected rotational velocity, and  derive accurate abundances, at least one high-resolution spectrum is needed for each star. To this end, an observational campaign is ongoing (Uytterhoeven et al. 2010a,b). Most stars of the \\dsct, \\gdor, and hybrid stars in our sample with magnitude Kp\\,$\\leq 10.5$\\,mag have recently been observed or are scheduled to be observed in the coming months. However,  70 \\% of the stars in Table\\,\\ref{classification} are fainter than magnitude Kp\\,$= 10.5$\\,mag, for which it is time-consuming and less practical to observe them with the available 2-m class telescopes that are equipped with a high-resolution spectrograph.  Because the oscillation modes in A-F type stars do not produce evident frequency patterns in their mode spectra, as is the case for solar-like oscillators, the identification of pulsation modes benefits from high-resolution spectral or multi-colour time series. Here we encounter limitations owing to the relative faintness of the {\\it Kepler} sample too. For instance, it is only feasible to efficiently spectroscopically monitor the few brightest (Kp\\,$\\leq 9$\\,mag) stars from the ground, while multi-colour photometry can go a few magnitudes fainter. Moreover, it will be impossible with the current instrumentation to detect the pulsation amplitudes of the order of a few $\\mu$mag from the ground. Therefore, only for a limited selection of the stars in Table\\,\\ref{classification}, i.e. bright stars exhibiting  high-amplitude variations, will it  be feasible to organize ground-based follow-up campaigns.  For all other stars, we will have to rely on extracting information on the pulsation modes directly from frequency patterns observed in the {\\it Kepler} data. Quasi-periodic patterns have been observed before in \\dsct\\,stars (Handler et al. 1997; Garc\\'{\\i}a Hern\\'andez et al. 2009). But in fast rotating stars the rotation destroys regular frequency and period patterns of p- and g-modes, which complicates the mode identification (e.g.\\,Ligni\\`eres et al. 2006; Ballot et al. 2010).  For slowly rotating  g-mode pulsators ($V_{\\rm rot} < 70$ km s$^{-1}$), a mode-identification technique has been developed that relies only on accurate values of at least three frequencies (Frequency Ratio Method, Moya et al. 2005; Su\\'arez et al. 2005), which is ideal to apply to the information extracted from the {\\it Kepler} white light, without colour or spectral information. Unfortunately, many of our stars are moderate-to-fast rotators (see Sect.\\,\\ref{sectChar}). Hence, the mode identification will be very challenging and will require  more investigation.  An individual analysis of the candidate hybrid stars is needed to confirm their hybrid status and to firmly characterize their pulsation properties. The current theoretical instability models for hybrid stars need to be revised to be able to accommodate all  stars that have been proposed as hybrid candidates in this paper. This includes a revision of the mechanisms that allow driving of p-\\,and g-modes in A-F type stars with a broad range of temperatures. Additional processes that can be investigated with possible effect on the driving are stochastic excitation (Houdek et al. 1999; Samadi et al. 2002), a convective driving mechanism similar to  g-mode pulsations in white dwarfs (Goldreich \\& Wu 1999),  a $\\kappa$\\,mechanism-related effect presented by Gautschy and L\\\"offler (1996) and L\\\"offler (2000), and radiative levitation (Turcotte et al. 2000). Asteroseismic diagnostics have been studied to find signatures of stochastic mechanisms at the origin of the instability of \\gdor\\,oscillators (Pereira et al. 2007). In that work, this possibility was not discarded, but continuous and precise space data were not yet available. The {\\it Kepler} time series of the sample of stars studied here will be an ideal new testbed for this method.  The long, continuous time series that {\\it Kepler} will deliver during its lifetime will unveil a large number of amplitudes at $\\mu$mag level. This precision will open up opportunities to search for signatures of granulation in the variable star light (Kallinger \\& Matthews 2010). Spectroscopically, convective signatures have been detected in the microturbulence and line broadening of A-F type stars cooler than \\Teff $= 10,000$\\,K (Landstreet et al. 2009).  Also the theoretical instability strips of the \\gdor\\,and \\dsct\\,pulsators need revision. As shown in Fig.\\,\\ref{Teffloggsample}, stars exhibiting purely \\gdor\\,or \\dsct\\,pulsations seem to exist beyond the current blue and red edge of the respective instability strips. Moreover, it is worth investigating if the evolutionary phase of \\gdor\\,stars can be derived from properties in their frequency spectra, as is recently suggested by Bouabid et al. (2011), based on a theoretical study of seismic properties of MS and pre-MS  \\gdor\\,pulsators.  Another open question is the existence of non-variable A-F type stars inside the instability strips. So far, it is suggested (Poretti et al. 2003; Breger 2004) that all seemingly constant stars in the instability strip are low-amplitude pulsators. In this study we find indications that non-variability exists within the instability strip, but a more in-depth investigation based on a more comprehensive sample of stars with precise values of \\Teff\\,and \\logg\\,is needed to confirm this.  Furthermore, candidate \\gdor\\,stars with only a few excited dominant modes deserve to be looked at in more detail. The relation between rotation and pulsations is not yet clear. Moreover, the differentiation between pulsations and rotational variability proves to be very difficult (Breger 2011; Monnier et al. 2010). In the pilot study by Balona et al. (2011d) it was suggested that pulsation and rotation periods might be very closely related. It needs to be investigated to which extent the rotation influences the excitation of the observed modes. To help this investigation, \\vsini~values are needed.  Constraints on important physical parameters that are crucial for seismic modelling, such as stellar radius and mass, can be derived directly for pulsators in binary systems (e.g. Tango et al. 2006; Desmet et al. 2010). Our sample consists of several binaries and eclipsing systems with (a) pulsating component(s) (see Table\\,\\ref{other}). Hence, these targets in particular are very promising for dedicated ground-based follow-up observations, and a seismic analysis. Moreover, it will be interesting to investigate the effect of tidal interactions on pulsation frequencies (e.g. Uytterhoeven et al. 2004; Derekas et al. 2011).  Four of the EBs with a candidate \\gdor, \\dsct, or hybrid component in our sample are known as chemically peculiar stars (see Table\\,\\ref{database}). Three candidate hybrid stars (out of 61 stars with known spectral type), four candidate \\dsct\\,stars (out of 67 stars), and one candidate \\gdor\\,star (out of 25 stars) are also Ap or Am stars. So far, we detected both p- and g-mode pulsators among the chemically peculiar stars. Balona et al. (2011c) stated that the instability strip of pulsating Am type stars and \\dsct\\,stars do not differ much. With the current small number statistics, it is not clear whether Ap/Am stars are indeed rare among \\gdor\\,stars (Handler \\& Shobbrook 2002). One of the open questions is if chemical peculiarity is related to hybridity. The first discovered hybrid HD\\,8801 (Henry \\& Fekel 2005) intruigingly turned out to be an Am star. In a recent abundance study by Hareter et al. (2011) one of the two studied hybrid stars is also confirmed as being a chemically peculiar star. Together with the results of this study,  this brings the total of known chemically peculiar hybrid stars to five. There is currently no evidence for a direct link between chemical peculiarity and hybrid behaviour, but a careful abundance analysis of a  representative sample of hybrid stars is needed to confirm this.  Many more (candidate) \\dsct, \\gdor, and hybrid stars are expected to be among the stars observed by {\\it Kepler}. Debosscher et al. (2011) reported the discovery of many additional \\dsct\\,and \\gdor\\,candidates in the public {\\it Kepler} Q1 data. Also, a considerable fraction of the host stars of the recently published 1235 {\\it Kepler} planet candidates (Borucki et al. 2011)  turn out to be A-F type stars.  Hence, we have promising prospects in studying and understanding the A-F type star variable behaviour in detail through a much larger and more complete sample of A-F stars in the {\\it Kepler} field when longer timestrings of {\\it Kepler} data will become publicly available.   {\\it Kepler} is definitely opening the window towards the accurate characterization of pulsating A-F stars."
    },
    "1107/1107.2106_arXiv.txt": {
        "abstract": "Spectroscopic follow-up of dozens of transiting planets has revealed the degree of alignment between the equators of stars and the orbits of the planets they host.  Here we determine a method, applicable to spotted stars, that can reveal the same information from the photometric discovery data, with no need for follow-up.   A spot model fit to the global light curve, parametrized by the spin orientation of the star, predicts when the planet will transit the spots. Observing several spot crossings during different transits then leads to constraints on the spin-orbit alignment.   In cases where stellar spots are small, the stellar inclination, $i_s$, and hence the true alignment, rather than just the sky projection, can be obtained.  This method has become possible with the advent of space telescopes such as \\emph{CoRoT} and \\emph{Kepler}, which photometrically monitor transiting planets over a nearly continuous, long time baseline.  We apply our method to CoRoT-2, and find the projected spin-orbit alignment angle, $\\lambda= 4.7^\\circ \\pm 12.3^\\circ$, in excellent agreement with a previous determination that employed the Rossiter-McLaughlin effect.  The large spots of the parent star, CoRoT-2, limit our precision on $i_s$: $84^\\circ \\pm 36^\\circ$, where $i_s < 90^\\circ ~(> 90^\\circ)$ indicates the rotation axis is tilted towards (away from) the line of sight. ",
        "introduction": "Transit observations of several systems, e.g., HD 209458 (Silva 2003), TrES-1 (Charbonneau et al. 2007), HD 189733 (Pont et al. 2007), and CoRoT-2 (Alonso et al. 2008, Silva-Valio et al. 2010), reveal anomalous flux rises during transit, most likely the result of the planet occulting a dark spot on the stellar surface.   Spots on the stellar surface complicate the estimation of planetary parameters from transit photometry and can lead to errors in planet size measurements if not accounted for (e.g., Czesla et al. 2009).  On the other hand, spots introduce structure into out-of-transit and in-transit observations, which may offer new opportunities to learn about the planet, its orbit, and the host star.  Silva-Valio (2008) pointed out the possibility of estimating the rotation period of transit host-stars, if an observer is able to catch the transiting planet occulting the same spot on two consecutive transits.  The key to the method is measuring the displacement of the \u201cspot perturbation\u201d along the transit chord and modeling the displacement as due to rotation of the star.  Applying the method to Hubble Space Telescope observations of HD 209458b,  and implicitly assuming alignment in the projected spin-orbit alignment angle ($\\lambda=0^\\circ$) and a stellar inclination of $i_s=90^{\\circ}$,  Silva-Valio (2008) constrained the rotation period of the star to be either 9.9 or 11.4 days.  Rotation period estimates from long-term photometric observations of out-of-transit variability in HD 209458 find a value consistent with the 11.4 day estimate.  Dittmann et al. (2009), in a similar analysis of TrES-1b, generalized the method by considering when $\\lambda \\neq 0^{\\circ}$ and $i_s \\neq 90^{\\circ}$, but found that this freedom introduced large uncertainty in the estimation of the rotation period.  We propose to reverse the chain of inference.  Because space-based transit monitoring missions, such as \\emph{Kepler} and \\emph{CoRoT}, observe each target over a nearly continuous, long time-baseline, we can robustly determine the rotation period and rotational phase of the occulted spot at the time of each transit by fitting a spot model to the \u201cglobal\u201d light curve.  By comparing the position of the spot perturbation along the transit chord with the spot's rotational phase, we can constrain $\\lambda$ and $i_s$.  Constraints on the projected spin-orbit alignment angle via Rossiter-McLaughlin (RM) measurements (e.g., Gaudi \\& Winn 2007) have yielded unique science.  For instance, it appears that misaligned systems are preferentially found around F-type stars (Schlaufman 2010) which may point to a connection between alignment and the larger convective atmospheres of cooler stars (Winn et al. 2010).  Meanwhile, attempts are being made to understand the relative importance of two (or more) modes of planetary migration, which bring planets from their formation zone into close-in orbits (Morton \\& Johnson 2010).  Planet-planet scattering followed by tidal circularization (Chatterjee et al. 2008) and Kozai oscillations followed by tidal circularization (Fabrycky \\& Tremaine 2007), both lead to misalignment, while migration through a gaseous disk should, as currently understood (Lin et al. 1996), lead to alignment.  In this paper, we describe our method to constrain the spin-orbit alignment and demonstrate an application to the CoRoT-2 system.  A similar concept has been developed by Sanchis-Ojeda et al. (2011) and was applied to WASP-4b, finding $\\lambda= -1^{+14}_{-12}$ degrees.  Our method and analysis was developed independently of Sanchis-Ojeda et al. (2011), whose analysis was reported while this manuscript was in preparation. In \\S 2, we discuss the \\emph{CoRoT} observations of CoRoT-2, measure the timing of in-transit spot-crossings, and estimate the timing uncertainty. In \\S 3, we describe a spot model to the global CoRoT-2 photometry, and our model for the in-transit spot timings.  We confront the spot timing model with the CoRoT-2 spot timing measurements. We conclude in \\S 4 with a discussion of our results and contrast our method with that of Sanchis-Ojeda et al. (2011). ",
        "conclusions": "In this paper, we have described a new method to determine the spin-orbit alignment of a transiting planet, and applied it to CoRoT-2b, finding $\\lambda = 4.7 \\pm 12.3^{\\circ}$.  This planet was also the subject of Rossiter-McLaughlin measurements by Bouchy et al. (2008), who found $\\lambda = 7.2 \\pm 4.5^{\\circ}$.  Compared to the results of Bouchy et al. (2008), our result is consistent, though less precise.  We also placed a weak constraint on the stellar inclination, $i_s = 84 \\pm 36^{\\circ}$.  Although this constraint is uninformative, we expect that applying our method to systems with starspots that are smaller than the planet will yield accurate determinations of the stellar inclination.  This possibility illustrates one advantage of our method over the RM method, which is entirely insensitive to differences in the stellar inclination.  Our method possesses other advantages over RM measurements.  RM measurements require costly radial velocity follow-up, while our method makes use of ``free'' data, in the sense that the data was already obtained for other purposes.  RM measurements are generally only possible for stars with relatively large projected rotational velocities, ruling out slow rotators, or relatively fast rotators that are nearly pole-on.  The RM method requires relatively large transit depths; the smallest planet that has succumbed to Rossiter-McLaughlin measurements is the Neptune-sized HAT-P-11b (Winn et al. 2010).  While our method also favors large transit depths, we expect that the photometric capability and data provided by the \\emph{Kepler} mission should allow for constraints on the spin-orbit angle of systems with planets smaller than HAT-P-11b.  The method described in this paper is similar, in concept, to that described by Sanchis-Ojeda  et al. (2011), and applies similar geometrical constraints.  In contrast to Sanchis-Ojeda et al. (2011), our method relies on global light curve measurements of the rotational period and phase of the occulted spot.  This has the disadvantage of requiring a continuous, long time-baseline of observations, which is generally only possible with space-based photometric monitoring, such as provided by the \\emph{CoRoT} and \\emph{Kepler} missions.  However, there is a great advantage to knowing where the spot is on the stellar surface.  For instance, only one transit of a planet over a spot is required to tell rough spin-orbit alignment.  If we know, from out-of-transit monitoring, that the spot is on the approaching side of the star (i.e., negative rotational phase), then a spot-crossing near the beginning of the transit indicates a probable prograde orbit, while a spot-crossing near the end of the transit indicates a probable retrograde orbit (with some exceptions for nearly pole-on stellar rotation).  To achieve good precision, of course, our method favors multiple spot-crossings at varying spot phases.  We gratefully acknowledge the helpful suggestions of our referee, A. Lanza."
    },
    "1107/1107.0932_arXiv.txt": {
        "abstract": "We present spatially-resolved spectroscopy of the galaxy cluster \\as1101, also known as S\\`{e}rsic 159-03,  with \\chandra, \\xmm\\ and \\rosat,  and investigate the presence of soft X-ray excess emission above the contribution from the hot intra-cluster medium. In earlier papers  we reported an extremely bright soft excess component that reached 100\\% of the thermal radiation in the R2 \\rosat\\ band (0.2-0.4 keV), using the HI column density measurement by Dickey and Lockman. In this paper we use the newer Leiden-Argentine-Bonn survey measurements of the HI column density towards \\as1101,  significantly lower than the previous value, and show that the soft excess emission in \\as1101\\ is now at the level of 10-20\\% of the hot gas emission, in line with those of a large sample of clusters analyzed by Bonamente et al. in 2002. The \\rosat\\ soft excess emission is detected  regardless of calibration uncertainties between \\chandra\\ and \\xmm. This new analysis of \\as1101 indicate that the 1/4~keV band emission is compatible with the presence of WHIM filaments connected to the cluster and extending outward into the intergalactic medium; the temperatures we find in this study are typically lower than the WHIM probed in other X-ray studies. We also show that the soft excess emission is compatible with a non-thermal origin as the inverse Compton scattering of relativistic electrons off the cosmic microwave background, with pressure less than 1\\% of the thermal electrons. ",
        "introduction": "Evidence for an excess of soft X-ray radiation in clusters above the contribution from the virial gas was initially discovered by \\citet{lieu1996a,lieu1996b} in the extreme ultra-violet ($h \\nu \\sim 0.1$ keV) with \\euve, and then confirmed with several other instruments, notably in the \\rosat\\ 1/4 keV band (see  \\citealt{bonamente2002} for the analysis of a sample of 38 clusters, and \\citealt{durret2008} for a recent review of the literature). The excess emission is usually modelled as an additional thermal component of lower temperature ($kT \\sim 10^6-10^7$~K) or as a non-thermal power law. A plausible scenario for the soft excess is thermal emission from  warm filaments seen in projection towards clusters \\citep[e.g.,][]{mittaz2004,bonamente2005}, or from low-entropy dense gas within the cluster \\citep{cheng2005}. Alternatively, a non-thermal power law can be used to model the excess, indicative of relativistic electrons that scatter the CMB photons and emit by Compton scattering \\citep[e.g.,][]{sarazin1998,lieu1999}. Other explanations are also viable, e.g. multi-temperature structure due to several mergers during the lifetime of a cluster \\citep{lehto2010}. Given the limited spectral resolution of the current CCD detector technology, it has not been possible to prove conclusively which additional model is a better fit to the data \\citep[see, e.g.,][]{nevalainen2003,bonamente2005,werner2007}.  Among the factors that affect the detection of the soft excess emission, the column density of absorbing HI plays a primary role. In recent years, the Leiden-Bonn-Argentine survey (LAB) of \\cite{kalberla2005} has effectively superseded the \\cite{dickey1990} measurements, and significant differences between the two surveys are occasionally present. Among the clusters in which the excess has been reported in the literature, \\as1101\\ is an outstanding case in that the \\cite{dickey1990} value of $N_H=1.80 \\times 10^{20}$ cm$^{-2}$ differs significantly from the LAB value of $N_H=1.15 \\times 10^{20}$ cm$^{-2}$. A lower $N_H$ value implies less Galactic absorption, and therefore more of the detected soft X-ray flux from the cluster originates as the soft tail of the hot, virial gas, reducing the need for an additional emission component.  The goal of this paper is to perform a spatially-resolved spectroscopic study of \\as1101\\ using the available data from \\chandra\\ and \\xmm, needed to measure the temperature of the hot gas, and \\rosat, needed to test whether at soft X-ray energies the radiation is consistent with the low-energy tail of the hot gas, or whether an excess of emission is present. This paper improves on the initial detection of the \\rosat\\ soft excess in \\as1101\\ by \\cite{bonamente2001a} in two ways: the hot gas is measured from \\chandra\\ and \\xmm\\ data that are much more sensitive to hot gas than the \\rosat\\ observations alone, and the use of the new LAB value for $N_H$. The presence of a soft component in the spectrum is investigated by testing whether the best-fit temperature of the hot gas has a bandpass dependence, as observed in certain clusters \\citep[for example,][ see Section~\\ref{sec:tprofiles}]{cavagnolo2008,lehto2010}, by  modelling the hot gas with a single-temperature model and a two-temperature model to account for cooler gas seen in projection, and determine whether the soft X-ray flux lies above the prediction (Section~\\ref{sec:hardband}-\\ref{sec:nh}), and by performing a fit to the whole X-ray band with the inclusion of a soft component (Section~\\ref{sec:2apec-soft}).  \\as1101, also known as S\\`{e}rsic~159-03, is located at  R.A.=23h13m58.5s, Dec.= -42d43m39, and the redshift is z=0.058; for a cosmology of $H_0=73$~km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, $\\Omega_{\\Lambda}=0.73$, the distance to the cluster is $D_A=220$ Mpc, and 1 arcmin corresponds to 64 kpc. The paper is structured as follows: in Section~\\ref{sec:data} we describe the data reduction methods, in Section~\\ref{sec:tprofiles} we determine the temperature profiles of the clusters from the various observations, in Section~\\ref{sec:excess} we determine the presence of a soft component above the hot ICM model, in Section~\\ref{sec:interpretation} we provide a discussion on the possible origin of the excess emission, and in Section~\\ref{sec:conclusions} we present our conclusions. Throughout the paper errors are quoted at the 1-$\\sigma$, or 68\\%  confidence level, unless otherwise stated. ",
        "conclusions": "\\label{sec:conclusions} In this paper we have presented spatially resolved X-ray spectroscopy of \\as1101\\ with \\chandra, \\xmm\\ and \\rosat. Our observations show that calibration uncertainties in the 0.7-2 keV X-ray band remain between the \\chandra\\ and \\xmm\\ detectors, with the \\xmm\\ whole-band best-fit temperatures typically lower than those measured by \\chandra, as also found by \\cite{nevalainen2010}.  Our analysis confirms the presence of soft excess emission in the \\rosat\\ 1/4 keV band, when the lower $N_H$ value of \\cite{kalberla2005} is used, and regardless of calibration uncertainties between \\chandra\\ and \\xmm\\ in the 0.7-2 keV band. The  revised soft excess fluxes that we derived show that the level of soft X-ray excess emission in all clusters studied to date does not typically exceed $\\sim$20\\% of the hot gas emission in regions that are within the virial radius. In fact, using the scaling relation $r_{500}=(0.796\\pm0.015)/(h E(z)) (T/5 \\text{keV}))^{(1.61\\pm0.11)/3}$~Mpc from \\cite{vikhlinin2006}, a $kT=3$~keV cluster like \\as1101\\ has $r_{500} \\simeq$ 1~Mpc, and the \\rosat\\ 1/4 keV band excess is certainly within 6~arcmin, or approximately $\\sim$400~kpc; a similar situation applies to other clusters studied with \\rosat\\ in \\cite{bonamente2002}. A possible exception is the Coma cluster, for which we have shown that there is excess emission out to $\\sim$5~Mpc \\citep{bonamente2009}, and for a cluster at $kT=8$~keV such as Coma, $r_{500} \\simeq$ 1.9~Mpc. The case of Coma may be exceptional in that it is one of the closest and brightest clusters, featuring  higher S/N observations  than any of the other clusters, and therefore it is possible that higher resolution soft X-ray observations near the virial radius may reveal that other clusters have soft X-ray halos that exceed the extent of the virialized gas.  The soft excess emission was interpreted as both thermal and non-thermal emission, and found that our data are consistent with both models for the excess. A possible origin for the soft excess is from WHIM filaments that are expected to converge towards massive galaxy clusters, and therefore become visible in X-rays even in the absence of metal lines. The best-fit temperature measured in this study are typically lower than the WHIM probed in other X-ray studies \\citep[e.g.,][]{werner2008}. In the case of a non-thermal origin, we find that the relativistic electrons responsible for the emission would feature a pressure of less than 1\\% of the thermal electrons. The upcoming launch of eRosita \\citep{predehl2010}, which improves on \\rosat's large effective area at 1/4 keV and field of view, will enable further tests of the presence of warm gas near clusters such as the ones we have performed for this paper with \\rosat."
    },
    "1107/1107.2937_arXiv.txt": {
        "abstract": "Pulsar timing arrays (PTAs) are expected to detect gravitational waves (GWs) from individual low--redshift ($z\\ltsim 1.5$) compact supermassive ($M\\gta 10^{9}\\Msol$) black hole (SMBH) binaries with orbital periods of $\\sim 0.1 - 10 \\yr$.  Identifying the electromagnetic (EM) counterparts of these sources would provide confirmation of putative direct detections of GWs, present a rare opportunity to study the environments of compact SMBH binaries, and could enable the use of these sources as standard sirens for cosmology.  Here we consider the feasibility of such an EM identification.  We show that because the host galaxies of resolved PTA sources are expected to be exceptionally massive and rare, it should be possible to find unique hosts of resolved sources out to $z\\approx 0.2$.  At higher redshifts, the PTA error boxes are larger, and may contain as many as $\\sim 100$ massive-galaxy interlopers.  The number of candidates, however, remains tractable for follow-up searches in upcoming wide-field EM surveys. We develop a toy model to characterize the dynamics and the thermal emission from a geometrically thin gaseous disc accreting onto a PTA-source SMBH binary. Our model predicts that at optical and infrared frequencies, the source should appear similar to a typical luminous active galactic nucleus (AGN). However, owing to the evacuation of the accretion flow by the binary's tidal torques, the source might have an unusually low soft X-ray luminosity and weak UV and broad optical emission lines, as compared to an AGN powered by a single SMBH with the same total mass. For sources at $z\\sim 1$, the decrement in the rest-frame UV should be observable as an extremely red optical color. These properties would make the PTA sources stand out among optically luminous AGN, and could allow their unique identification. Our results also suggest that accreting compact SMBH binaries may be included among the observed population of optically bright, X-ray-dim AGN. ",
        "introduction": "\\label{sec:intro}  Over the last several years, the possibility of observing both the gravitational-wave (GW) and electromagnetic (EM) emission signatures of coalescing supermassive black hole (SMBH) binaries has received intense attention (\\citealt{HH05, Kocsis+06, Kocsis+07, Kocsis+08, Dotti+06}; for specific proposed mechanisms for EM signatures, see \\citealt{AN02} and \\citealt{MP05}, as well as recent reviews by \\citealt{Haiman+09} and \\citealt{Schnittman11}). The bursts of GWs emitted by such systems can now be predicted by numerical general relativity \\citep{Pretorius05, Baker+06, Campan+06}, and are expected to be observed by current and future detectors. The temporal evolution of the gravitational waveform can be used to extract the luminosity distance, help constrain the location of the source on the sky, and determine the masses and spins of the SMBHs.  If an EM signature of the coalescence can also be identified, this would allow for a determination of the source redshift, turning merging black holes into ``standard sirens'' for probing cosmic expansion.\\footnote{ The importance of such GW+EM observations for cosmography was first discussed by \\cite{Schutz86} in the context of merging neutron star binaries.}  Such multi-messenger observations would also enable astronomical investigations of SMBHs whose masses, spins and orbital parameters are already known, presenting ideal laboratories for investigating accretion physics in active galactic nuclei (AGN).  Furthermore, if major mergers of galaxies trigger luminous AGN activity (e.g., \\citealt{Sanders+88, Hernquist89, Carlberg90, BarnesHernquist91, HernquistMihos95, MihosHernquist96, KauffmannHaehnelt00, Hopkins+07a, Hopkins+08}), then the characteristic EM emission promptly following the SMBH coalescence may herald the birth of a quasar \\citep{THM10}.  To date, theoretical studies of EM signatures of GW-emitting SMBH binaries have largely centred on systems predicted to be detectable by future space-based laser interferometers such as {\\it Laser Interferometer Space Antennae} ({\\it LISA}), with total mass $\\sim 10^{5-7}(1+z)^{-1}\\Msol$ and redshifts as high as $z\\gta 10$. For the purposes of multi-messenger astronomy, a paramount feature of interferometer-detectors is the precision with they could determine the angular sky position of SMBH sources \\citep{Kocsis+06}: to $\\ltsim 1 \\deg$ \\citep{Vecchio04, LH08}, or perhaps even to $\\ltsim 1^{\\prime}$ when spin-induced precession \\citep{LH06} or higher-order harmonics \\citep{McWill+10} are included in the analysis of the waveform.  In this paper, we evaluate the prospects of performing multi-messenger using pulsar timing arrays (PTAs), which aim to detect low-frequency (nHz) GWs through precision timing of Galactic millisecond pulsars. A major goal for PTAs is to detect the stochastic GW background due to the population of compact SMBH binaries in our cosmic neighborhood ($z\\ltsim 2$). \\cite{Jenet+05} showed that this requires the timing of some $\\ga 20$ pulsars to $\\sim 100 ~{\\rm ns}$ precision over 5 years, a goal that could be achieved by currently operational arrays \\citep[][and refs. therein]{Manchester08, Verbiest+09}. Recent theoretical population-synthesis studies \\citep{SVV09, SesVec10, KocSes11} have predicted that PTA observations may also be able to {\\it individually resolve} the most massive and/or most nearby SMBH binaries that stick out above the stochastic background.\\footnote{In principle, PTAs can also detect GW-memory bursts from individual major SMBH mergers \\citep{Seto09, Pshirkov+10,HL10}. However, the number of detections is expected to be quite low, perhaps much lower than unity over a full observation campaign \\citep[see last paragraph in][]{Seto09}. We therefore limit the present paper to the population and EM signatures of compact SMBH binaries.} The primary factor that determines the number of detectable individual sources is expected to be the array sensitivity, with theoretical uncertainties such as the binary mass function and orbital evolution affecting predictions of the observable population by factors of unity \\citep{Sesana+11}.   The systems individually detected by PTAs will differ from {\\it LISA} targets in several ways. First, according to the population synthesis studies, the binaries will have chirp masses \\beq \\Mch\\equiv M_{1}^{3/5}M_{2}^{3/5}M^{-1/5}=\\eta^{3/5}M, \\label{eq:Mch} \\eeq typically around $\\Mch\\sim10^{8.5}\\Msol$. Above, $M_{2}\\le M_{1}$ are the masses of each member of the binary, and $\\eta\\equiv M_{1}M_{2}/M^{2}\\le 1/4$ is the symmetric mass ratio. In other words, the systems of interest here are much more massive than those relevant for {\\it LISA}. Second, PTA sources lie at much closer cosmological distances. While the redshift probability distribution is poorly known, it is expected to decline steeply outside the range $0.1\\ltsim z\\ltsim 1.5$. The low cutoff is due to the smaller volume contained within $z$, whereas the high cutoff is due to the attenuation of the GW signal and the decline of the intrinsic SMBH merger rate. Third, the PTA targets are not as compact, having periods of $P\\sim 1\\yr$. The larger separations correspond to slower orbital decay; the vast majority of PTA sources will not coalesce within a human lifetime. That the orbit --- and thus the waveform --- is expected to evolve much less appreciably with time poses a particularly difficult challenge for determining the masses and luminosity distances of PTA sources. Particularly important for the prospects of performing synergistic EM observations on these sources is their relatively poor sky-localization. The solid-angle uncertainty in the source position may be anywhere from $\\Delta \\Omega \\ltsim 3\\deg^{2}$ (in the best-case scenario in which the contributions to the signal from individual pulsars are known; \\citealt{CorCor10}, hereafter CC10) to as large as $\\Delta \\Omega \\sim 40\\deg^{2}$ (if individual pulsar contributions cannot be extracted from the data; \\citealt{SesVec10}, hereafter SV10).  Several tell-tale EM features of such compact SMBH binaries have been proposed in the literature, including periodic emission modulated at the orbital frequency of the PTA source (e.g., \\citealt{Haiman+09} and references therein) and double-peaked broad emission lines \\citep[e.g.,][]{Gaskell96, Zhou+04, Bogdan+08, BlechaLoeb08, BorLau09, Bogdan+09b, Dotti+09,TanGri09, ShenLoeb10}. In this work, we will explore in detail the suggestion by \\cite{MP05} that tidal torques of a near-merger SMBH binary will have evacuated the central region of its accretion disc, resulting in a markedly softer thermal spectrum.  Below, we investigate whether individually resolved PTA sources may be viable targets for EM identification. In particular, we address the following two questions:  \\begin{enumerate}  \\item What is the average number $N_g$ of candidate host galaxies --- that is, interloping galaxies that could plausibly harbour the GW source --- in a typical error box of a PTA detection?  Of particular interest is whether there are plausible scenarios for detecting individually resolved sources with $N_g<1$, i.e., cases where the source may be uniquely identified with an EM search of the three--dimensional PTA error box.  \\item In cases where $N_g>1$, what can be done to distinguish the true host galaxy of the source from the other interlopers?  Motivated by the hypothesis that galaxy mergers can fuel AGN activity, we will consider the differences in thermal emission properties predicted by disc models of AGN powered by a compact SMBH binary as opposed to one powered by a solitary SMBH of the same total mass.  \\end{enumerate}    This paper is organized as follows.  In \\S \\ref{sec:hosts}, we provide a brief overview of the expected population of PTA-resolved SMBH binaries as well as the anticipated detection error box of such objects.  We consider the error box prescriptions of \\citetalias{SesVec10} and \\citetalias{CorCor10}, characterize the types of astronomical objects that are plausible hosts of a PTA-resolved binary, and estimate the number of such objects.  In \\S 3, we describe a toy model to calculate the dynamical state and thermal emission features of gas accreting onto a resolved SMBH binary.  We discuss several features predicted by the model which, if observed, could help the EM identification of individually resolved PTA sources. We summarize our findings and offer our conclusions in \\S 4.  Throughout this paper, $c$ denotes the speed of light; $G$ is the gravitational constant; $k_{\\rm B}$ is the Boltzmann constant; $\\sigma_{\\rm SB}$ is the Stefan-Boltzmann constant; and $m_{\\rm p}$ is the mass of the proton. ",
        "conclusions": "In this paper we considered the possibility that individually resolved PTA sources --- SMBH binaries with $M\\sim 10^{9}\\Msol$, $M_{2}/M_{1}\\sim 1/4$, $P\\sim 0.1-1 \\yr$ and $z\\ltsim 1.5$ --- may be identified with EM observations if they reside in gas-rich environments.  Multi-wavelength observations of such systems would allow for studies of an AGN in a system that is known to harbour a compact SMBH binary, thus providing a unique window into gas accretion in a rapidly time-varying gravitational potential.  If, as suggested by \\cite{CorCor10}, PTAs can constrain the luminosity distance to individually resolved sources, these SMBHs can be used as ``standard sirens'' to measure the cosmic expansion history.  Interestingly, the predicted redshift distribution of these sources lies between $0.1 \\ltsim z \\ltsim 1.5$, a range comparable to that of the deepest Type-Ia supernovae surveys and where {\\it LISA} detection rates are expected to be low \\citep{Sesana+07}.  Our findings can be summarized as follows: \\begin{itemize}  \\item The number of interloping massive halos and AGN that may be confused with the PTA source is typically $N_g\\sim 10^{4}$ for a $10^{9}\\Msol$ binary if the contributions to the signal from individual pulsars are not identifiable in the PTA data.  \\item In the more optimistic case, the pulsar term can be constrained and utilized to better determine the sky location of the source as well as constrain its redshift and mass. The number of interlopers can then drop to $N_g\\sim 10-100$ at $z\\sim 0.5$, and perhaps to $N<1$ for close sources at redshifts as close as $z\\ltsim 0.2$.  \\item By considering the orbital evolution history of an accreting PTA source, first by tidal interactions with the circumbinary gas and then by GW emission, we showed that a gaseous accretion disc around the source can be expected to be gas-poor both inside and immediately outside its orbit.  They would thus have optical and infrared luminosities comparable with typical quasars, while exhibiting low thermal X-ray luminosities and weak UV emission lines. The downturn in flux below $300~{\\rm nm}$ could be discerned by optical observations if the source redshift is $z\\ga 1$. The leakage of circumbinary gas into the cavity may shock the circumprimary (-secondary) disc and cause substantial quasiperiodic fluctuations in the UV and X-ray. Searching for AGN in the PTA error box with one or more of these atypical characteristics could lead to the identification of a single EM counterpart.  Further monitoring candidate counterparts for periodicity and other theoretically predicted pre-coalescence signatures may also aid identification. \\item The above emission features are general for AGN powered by thin accretion discs around compact SMBH binaries. They could thus be used to discover such systems even in the absence of a GW detection. This is a particularly interesting possibility, as a small fraction ($\\ltsim 1\\%$) of optically luminous AGN is known to have low X-ray luminosities and unusually weak emission lines. Upcoming wide-field surveys in the optical and X-rays should together discover many more such AGN, as well as observe them over temporal separations of weeks to months. Such cadential, multi-wavelength data may lead to the first observations of compact SMBH binaries with sub-parsec separations. \\end{itemize}  We have assumed that a PTA source would appear as luminous AGN, with the gaseous fuel perhaps being supplied by the preceding merger of the binary's host galaxies. The degree to which galaxy mergers dictate AGN activity remains an open question, and the link between SMBH binarity and AGN activity is even less certain. It is possible that many SMBH binaries that are individually detected by PTAs will have no EM counterpart at all.  Our results were calculated using a simple semi-analytic accretion disc model, a central assumption of which is that the binary's tidal torques are able to open a central cavity in the disc. In the radiation-dominated regions of interest, strong horizontal advective fluxes or vertical thickening of the disc may act to close such a cavity and wipe out the features we predict. The features would also not be present if the disc and binary's orbits do not lie on the same plane, as in the binary model of the variable BL Lac object OJ 287 \\citep{LV96}. Absorption and reprocessing by the binary's host galaxy may also act to mask or mimic the intrinsic thermal AGN emission we have modeled."
    },
    "1107/1107.4793_arXiv.txt": {
        "abstract": "We carry out a spectral analysis of the archival {\\it FUSE} spectrum of the VY Scl nova-like  cataclysmic variable MV Lyrae obtained in the high state. We find  that standard disk models fail to fit the flux in the shorter wavelengths of {\\it FUSE} ($\\lambda < 950$\\AA ). An improved  fit is obtained by including a modeling of the boundary layer at the inner edge of the disk. The result of the modeling shows that in the high state the disk has a moderate accretion rate $\\dot{M} \\approx 2 \\times 10^{-9} M_{\\odot}$/yr, a low inclination, a boundary layer with a temperature of $\\sim 100,000$K, and size $\\sim 0.20R_{wd}$, and the white dwarf is possibly heated up to a temperature $T \\approx 50,000$K or higher. ",
        "introduction": "\\subsection{The VY Scl Nova-like System MV Lyrae}  Cataclysmic variables (CVs) are evolved compact binaries in which the primary, a white dwarf (WD), accretes matter and angular momentum from the secondary, a star approximately on the main sequence filling its Roche-lobe. In dwarf nova (DN) and non-magnetic nova-like (NL) systems (two classes of CVs), the matter is transferred by means of an accretion disk around the          WD. DN systems are found mostly in a faint, quiescent state with a low mass transfer rate ($\\dot{M}\\sim 10^{-12}-10^{-11}M_{\\odot}$/yr), and every few weeks to months they transit into a bright, outburst state lasting days to weeks, as the mass transfer suddenly increases by several orders of magnitude ($\\dot{M}\\sim 10^{-9}-10^{-8}M_{\\odot}$/yr). The transition to outburst in DNs is believed to be due to a disk instability \\citep{sha86,can88,can98}. NL systems are found mainly in a high state phase, characterized by a large mass accretion rate, and from time to time the mass accretion rate suddenly drops by several orders of magnitude. The transition from high state to low state in NL is caused by a throttling of the mass transfer, possibly due to magnetic activity of the secondary star (see \\citet{war95} for a review of CVs). With time, the accreted matter in CV systems forms a hydrogen-rich layer that covers the surface of the          WD. When enough material is accreted (every few thousand years), the hydrogen-rich material undergoes a thermonuclear runaway and the system forms a classical nova event.   MV Lyr is a VY Scl nova-like system, spending most of its time in the high state and occasionally undergoing short-duration drops in brightness. When this happens, the magnitude of MV Lyr drops from $V\\approx 12-13$ to $V\\approx 16-18$ \\citep{hoa04}. From the AAVSO online data (http://www.aavso.org/) it appears that MV Lyr can stay in the high state for up to $\\sim 5$yr, and then for a period of a few years to $\\sim10$yr it spends its time alternating between high state and low state on a time scale of a few months to a year. The orbital period of the binary in MV Lyr is $P=3.19$hr and the mass ratio was found to be $q = M_{2nd}/M_{wd} =0.4$ \\citep{sch81,ski95}. Because of the small radial velocity amplitudes and the lack of eclipses, the inclination is constrained to be in the range $i=10^{\\circ}-13^{\\circ}$, though it has been remarked \\citep{lin05} that the inclination could be as low as $i=7^{\\circ}$.  Observations of MV Lyr during a low state with the {\\it International Ultraviolet Explorer ({\\it IUE})} indicated a hot          WD possibly reaching 50,000K \\citep{szk82} or higher \\citep{chi82}. These observations were later confirmed when MV Lyr was observed with {\\it FUSE} on July 7, 2002, during a low state that lasted for about 8 months. \\citet{hoa04} carried out a spectral analysis of the {\\it FUSE} spectrum and found that the          WD must have a temperature of 47,000K, a gravity $\\log{g}=8.25$, a projected rotational velocity $V_{rot} \\sin{i} =200$km/s, subsolar abundances $Z = 0.3 \\times Z_{\\odot}$, and a distance of $505 \\pm 50$pc. At the time of the observations, the magnitude of MV Lyr was possibly as low as $V\\approx 18$ with a mass accretion rate no more than $\\dot{M} \\approx 3 \\times 10^{-13}M_{\\odot}$/yr \\citep{hoa04}. A follow up analysis \\citep{lin05} was carried out to study the different states of MV Lyr and its secondary, based  on archival {\\it IUE} spectra, a HST/{\\it STIS} snapshot and optical spectra. Using the parameters they derived in \\citet{hoa04}, \\citet{lin05} found that standard disk models did not provide adequate fits to the spectra; instead, the disk had to be modified by, for example, a truncation of the inner disk, an isothermal radial temperature profile in the outer disk, or both. These modified disk models provided some satisfactory fit to the spectra obtained in the high state, leading to a mass accretion rate of the order of $3 \\times 10^{-9}M_{\\odot}$/yr.  In the present analysis, the standard disk model fails again to provide a good fit to the observed ({\\it{FUSE}}) spectrum of the system in the high state. This, however, points to a more general problem, namely that the standard disk model has failed to fit the observed spectra of many other CVs observed in the high state \\citep{pue07}. The direct  implication is that the standard disk model itself is inadequate in the ultraviolet (UV). It is in this context that we analyze here the {\\it FUSE} spectrum of MV Lyr in its high state. For this reason, we review next the problems encountered in the modeling of CVs in the high state.  \\subsection{UV Spectral Modeling of CVs in the High State}  In high state systems, the UV emission predominantly originates from the disk,  and for that reason these systems became of special interest, as they were expected to have an accretion disk radial temperature profile given by the analytical expression $T_{eff}(r)$ of the standard disk model. More than two decades ago, \\citet{wad88} used {\\it IUE} spectra to show that CV NL systems (high state systems)  systematically disagree with the standard disk model when either blackbody spectra or Kurucz stellar model spectra were used to represent the accretion disk. A much improved version was developed by \\citet{hub90} for modeling annuli of accretion disks that explicitly included calculation of synthetic spectra using the standard disk model (the codes TLDISK/TLUSTY, SYNSPEC and DISKSYN), and spectra of disks were computed \\citep{wad98} assuming the standard disk model \\citep{sha73,lyn74,pri81}. UV spectra (HST, HUT, {\\it IUE, FUSE}) of DNs in outburst and high brightness states of NLs were successfully modeled for a number of systems with synthetic spectra of optically thick disks \\citep{kni02,ham07}.  However, with time, it became clear that a significant number of systems still could not  be fitted using the standard disk model \\citep{lon94,kni97,rob99} as first shown by \\citet{wad88}. \\citet{oro03} suggested a disk temperature profile ($T(r)$) flatter than the standard disk model, to remediate at least partially to the problem, invoking a possible non-steady mass accretion rate and/or irradiation in the outer part of the disk. More recently, the UV spectra of RW Sex, V3885 Sgr, \\& UX UMa \\citep{lin08,lin09,lin10} were also modeled using a modified temperature profile, and it was suggested  that the outer disk temperature might increase due to its interaction with the impacting L1 stream. Maybe the most significant of all is the statistical study of \\citet{pue07} of 33 disk-systems which indicates that a revision of the standard model temperature profile, at least in the innermost part of the disk, is required.  These studies concentrated  mainly on using          UV spectra (e.g. {\\it IUE}, HST/{\\it STIS}, {\\it FUSE}), and the contribution to the UV continuum comes mainly from regions in the system where the temperature ranges roughly between $\\sim$10,000K and $\\sim$100,000K. Lower temperatures peak in the optical, while higher temperatures peak in the extreme UV. The          WD star and the disk are the main contributors to the UV flux, while  the secondary star and the bright spot (the edge of the disk impacted by the incoming stream of matter) peak in the optical and near UV, respectively. For disks in CVs with mass accretion rate $\\sim 10^{-8}M_{\\odot}$/yr, the inner disk reaches a temperature higher than 50,000K. These disk models, therefore, have a significant UV flux contribution from the inner region of the disk.  That same innermost part of the disk is the region where the matter in the disk has to decelerate from its Keplerian motion in order to be accreted onto the more slowly rotating stellar surface, as the centrifugal force keeps the matter from accreting. This region is known as the boundary layer (BL) and up to half the accretion luminosity can be released in this final phase of the accretion process. Therefore, the emission and heating of the BL have to be taken into account when modeling an accretion disk around a non-  (or weakly- ) magnetic star \\citep{lyn74,pri81}. It has now become clear that the temperature profile in the inner disk is affected by the heating from the BL. Most importantly, BLs in high state have a temperature of the order of 100,000K-200,000K \\citep{pop95,god95,pir04}, which is much lower than the first analytical estimates of $\\sim$500,000K \\citep{pri77,pri79}. Consequently, the BL of CV disk systems in the high state contributes significantly to the UV, even though it was predicted to emit only in the soft X-ray range (see also next section). Therefore, the BL has to be included in the disk modeling and the inner disk temperature profile has to be modified accordingly.    It is the purpose of the present work to include an elementary (and preliminary) model of the BL in the spectral modeling of the {\\it FUSE} spectrum of MV Lyr taken during its high state. Since the BL has an elevated temperature, its contribution will be more pronounced in the shorter wavelengths of the spectral range of {\\it FUSE}.  In the next section a brief review of the standard disk and BL models are given, and it is shown that the theory does indeed predict a BL contributing a non-negligible flux in the far          UV. The {\\it FUSE} spectrum of MV Lyr together with the details of the synthetic spectral modeling are presented in section 3. The results follow in section 4, and are discussed in the concluding section. ",
        "conclusions": "Overall, the best fit to the {\\it FUSE} spectrum of MV Lyr in its high state consists of a composite WD + disk + BL model with a mass accretion rate  $\\dot{M} = 2 \\times 10^{-9}M_{\\odot}$/yr. There are two equally acceptable BL solutions, either a broad extended (two-ring) BL with a rather low temperature (in the range 100,000K-150,000K), or a thin (one-ring) BL with a higher temperature (175,000K). The WD temperature is probably higher than the 44,000K found in the low state, due to ongoing accretion but it affects only slightly  the quality of the fitting of the models, we discuss this further at the end of this section.  Since the modeling of the {\\it FUSE} spectrum generates more than one solution over a finite spectral range, one must consider the bolometric luminosity of each component for comparison with equations (2) \\& (6). For each model, the total luminosity of the BL is computed assuming it radiates as a black body. This is a reasonable assumption, as in this regime the BL is most probably optically thick \\citep{god95,pop95}. The two components of temperature $T_1$ \\& $T_2$ (given in Table 2), and with radii $r_1=1.05R_*$ and $r_2=1.2R_*$ are added. One then simply uses the Stefan-Boltzmann law and area of each ring to find the total        BL      luminosity. The bolometric luminosity obtained is then compared with the theoretical expectation in eq.(6) using eq.(2). For almost all the disk+BL models listed in Table 2 it is found that the computed bolometric BL luminosity is larger than the disk luminosity. The discrepancy is even larger for the model with $\\dot{M} = 1 \\times 10^{-9}M_{\\odot}$/yr. The only models for which the BL bolometric luminosity is smaller than the disk luminosity (eq.6) is for an extended BL (two-ring model) with a temperature as low as 100,000K. For that model the bolometric BL luminosity is 93\\% of the disk luminosity, implying a rotational velocity (eq.6) of less than 4\\% its Keplerian value or about 135km/s.   With an inclination $i=10^{\\circ}\\pm 3^{\\circ}$ and a projected rotational velocity of $200\\pm 50$km/s (as derived in the low state), the (non-projected) rotational velocity should be in the range 667km/s$< V_{rot}<$2050km/s, or $0.18 < \\Omega_*/\\Omega_K < 0.54$ (the lower limit is for a velocity of 150km/s and $i=13^{\\circ}$, while the upper limit is for 250km/s and $i=7^{\\circ}$).  This produces a BL luminosity in the range $ 0.68 > L_{BL}/L_{disk} > 0.21$. For the BL model to agree with the upper limit, one could reduce the BL temperature (by 10,000K)  or increase the mass accretion rate (by 1/3, such a model is listed at the end of Table 2) or both. However, one notes that the error in the WD mass of $0.1M_{\\odot}$ ($\\sim 0.2$ in $\\log{g}$)  produces a computed $L_{BL}/L_{disk}$ between 0.55 (for a $0.9M_{\\odot}$ WD mass) and $\\sim$1.5 (for a $0.7M_{\\odot}$ WD mass).  Since the BL model used here is rather simplistic, a perfect agreement between the computed bolometric luminosity of the BL and that computed from eq.6 is not expected. However, within the limits of the errors, the results of the BL model are consistent with a broad BL (of size $\\sim 0.2 R_*$) with a rather low temperature ($\\sim 100,000$K or slightly less) and a mass accretion rate of the order of $ \\approx 2 \\times 10^{-9} M_{\\odot}$/yr (or slightly larger). These models are located the lower part of Table 2 and the fit improves (lower $\\chi^2_{\\nu}$ and better distance) as the temperature of the WD increases, reaching a best fit for $T=70,000$K. Though the temperature of the WD is certainly not that high, this certainly points to the fact that the temperature of the WD is elevated during the high state reaching $\\sim 50,000$K or a little higher."
    },
    "1107/1107.5276_arXiv.txt": {
        "abstract": "We used the DEIMOS spectrograph on the Keck II Telescope to obtain spectra of galaxies in the fields of five distant, rich galaxy clusters over the redshift range $0.5 < z < 0.9$ in a search for luminous, compact, blue galaxies (LCBGs).  Unlike traditional studies of galaxy clusters, we preferentially targeted blue cluster members identified via multi-band photometric pre-selection based on imaging data from the WIYN telescope.  Of the 1288 sources that we targeted, we determined secure spectroscopic redshifts for 848 sources, yielding a total success rate of $66\\%$.  Our redshift measurements are in good agreement with those previously reported in the literature, except for 11 targets which we believe were previously in error.  Within our sample, we confirm the presence of 53 LCBGs in the five galaxy clusters. The clusters all stand out as distinct peaks in the redshift distribution of LCBGs with the average number density of LCBGs ranging from $1.65\\pm0.25~\\mbox{Mpc}^{-3}$ at $z=0.55$ to $3.13\\pm0.65~\\mbox{Mpc}^{-3}$ at $z=0.8$.  The number density of LCBGs in clustes exceeds the field desnity by a factor of $749\\pm116$ at $z=0.55$; at $z=0.8$, the corresponding ratio is $E=416\\pm95$.  At $z=0.55$, this enhancement is well above that seen for blue galaxies or the overall cluster population, indicating that LCBGs are preferentially triggered in high-density environments at intermediate redshifts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.0665_arXiv.txt": {
        "abstract": "Very energetic cosmic rays entering the atmosphere of the Earth will create a plasma cloud moving with almost the speed of light. The magnetic field of the Earth induces an electric current in this cloud which is responsible for the emission of coherent electromagnetic radiation. We propose to search for a new effect: due to the index of refraction of air this radiation is collimated in a Cherenkov cone. To express the difference from usual Cherenkov radiation, i.e.\\ the emission from a fast moving electric charge, we call this magnetically-induced Cherenkov radiation. We indicate its signature and possible experimental verification. ",
        "introduction": "It is a well-known fact that an emitting source moving at a velocity exceeding the wave-propagation velocity in the medium will induce a shock-wave. Prime examples of this phenomenon are the sonic boom emitted by a super-sonic airplane, a bow-wave from a moving ship, and Cherenkov radiation emitted by an electric charge moving at almost the vacuum speed of light in a medium with an appreciable index of refraction. In this work we show that a similar effect occurs  in the emission from a fast moving electric current. It is suggested that this effect manifests itself in the emission of electromagnetic waves from particle cascades in the atmosphere of the Earth initiated by ultra-high energy (UHE) cosmic rays, with energies in excess of $10^{17}$\\,eV.  An UHE cosmic ray entering the atmosphere of the Earth creates a cascade of particles, called an extensive air shower (EAS). In this cascade there are copious amounts ($>10^6$, depending on the initial energy of the cosmic ray) of electrons and positrons forming a small plasma cloud. This cloud, with a typical size of less than 1\\,m, moves with almost the light velocity towards Earth. The magnetic field of the Earth, by means of the Lorentz force, induces a drift velocity for the leptons which is perpendicular to the direction of the initial cosmic ray and opposite for electrons and positrons. As a result an electric current is created in the fast moving plasma cloud. The strength of the induced electric current is roughly proportional to the number of charged particles in the plasma cloud as the induced drift velocity is varying little with height. This macroscopic picture~\\cite{Sch08} was recently confirmed~\\cite{Hue10-Arena} to agree with a microscopic description~\\cite{Lud10-Arena}. Even when the index of refraction of air would be equal to that of vacuum this varying electric current emits electromagnetic waves and coherent emission occurs at a wavelength longer than the size of the charge cloud, i.e.\\ for radio frequencies $\\nu< 300$\\,MHz~\\cite{Sch09-Arena}. The geomagnetic emission mechanism~\\cite{Kah66} has been confirmed from data~\\cite{Fal05}. In addition to the induced current the plasma cloud has a net charge excess which also radiates. The polarization direction of radiation distinguishes the emission due to the charge excess and the geomagnetic current~\\cite{dVries10,Schoorl11}. Since charge-excess radiation is generally smaller in intensity we will concentrate in this work on geomagnetic emission.  As is well known the propagation speed of electromagnetic waves is $c/n$ where $c$ denotes the velocity of light in vacuum. In this work we investigate the effect of the index of refraction of air, $n$, on the emission following Ref.~\\cite{Wer08}. The effects of Cherenkov radiation from EAS have also been addressed in Ref.~\\cite{Kalmykov2006347}. In this work we show that for realistic values for $n$ the Cherenkov effect introduces distinct features in the ground pattern of the emitted radiation. ",
        "conclusions": ""
    },
    "1107/1107.5040_arXiv.txt": {
        "abstract": "\\noindent Accurate relative spectrophotometry is critical for many science applications. Small wavelength scale residuals in the flux calibration can significantly impact the measurements of weak emission and absorption features in the spectra. Using Sloan Digital Sky Survey data, we demonstrate that the average spectra of carefully selected red-sequence galaxies can be used as a spectroscopic standard to improve the relative spectrophotometry precision to 0.1\\% on small wavelength scales (from a few to hundreds of Angstroms). We achieve this precision by comparing stacked spectra across tiny redshift intervals. The redshift intervals must be small enough that any systematic stellar population evolution is minimized and less than the spectrophotometric uncertainty. This purely empirical technique does not require any theoretical knowledge of true galaxy spectra. It can be applied to all large spectroscopic galaxy redshift surveys that sample a large number of galaxies in a uniform population.  \\rightskip=0pt ",
        "introduction": "In astronomical spectroscopy, accurate spectrophotometric calibration is critical for many science applications, but difficult to achieve. All calibrations require a standard source whose intrinsic properties we know. By observing it with the same system (atmosphere, telescope, instrument, detector) as used for observing the science targets, we can infer the throughput of the system as a function of wavelength. However, the accuracy of this calibration is limited by how well we know the standards and their measurement uncertainty. For ground-based astronomical observations, we have to rely on natural standards (but see \\citealt{Kaiser08,Kaiser10} and \\citealt{Albert11} for attempts to use satellite-mounted standards). Nearly all extragalactic spectroscopy has used stars as the calibration standards \\citep[e.g.,][]{Stoughton02}. Therefore, the accuracy is limited by how well we know the true spectra of stars and the uncertainty in the standard star spectra.        The current state-of-the-art spectrophotometric calibration utilized in large surveys can be found in the Sloan Digital Sky Survey (SDSS; \\citealt{York00}; \\citealt{SDSSDR7}). In SDSS, each plate includes a set of 16 standard stars, which are color-selected to be F8 subdwarfs. Their spectra are reduced in the same way as science targets. The calibration is achieved by comparing the galactic-reddening-corrected spectra of these stars to a grid of theoretical spectra generated from Kurucz model atmospheres. The average ratio of each star to its best-fit model is taken as the flux calibration vector and applied to all the science targets \\citep{SDSSDR2}. Finally, the corrected spectra are put on an absolute scale by comparing the synthetic photometry to fiber magnitudes from the SDSS photometry.  The accuracy of this method is limited by the signal to noise of the standard star spectra and the accuracy of the models. The former limits the confidence in using them to correct for the system response variation on small scales. In SDSS, a fourth-order b-spline with 50 breakpoints is used to smooth the flux calibration vector before applying it to calibrate science targets, thus leaving small-scale variations uncorrected. The theoretical model of the standard stars could also have unknown systematic residuals on small scales.  Overall, the wavelength-dependent relative flux calibration in SDSS is good to the level of 1\\%-2\\%. However, this accuracy is still inadequate in certain science applications. % The problem is illustrated in Figure~\\ref{fig:diagnosis_before}. Here, we show the \\nii/\\sii\\ flux ratio and $D_n(4000)$ index for a population of galaxies that have emission lines with ratios characteristic of low-ionization nuclear emission-line regions (LINERs). The sample is volume-limited. Their \\nii/\\sii\\ ratios vary systematically with redshift, at a level nearly comparable to the intrinsic scatter at a single redshift. Similarly, the $D_n(4000)$ index also show unphysical variations with redshift.  These emission line fluxes and spectral indices are measured with our own software, which is an improved version of the code described by \\cite{Yan06}. Similar systematics also exist if we adopt the DR7 catalog produced by the MPA--JHU collaboration\\footnote{http://www.mpa-garching.mpg.de/SDSS/} \\citep{TremontiHK04}, as shown in the right panels of Figure~\\ref{fig:diagnosis_before}. The detailed wiggles differ between our measurements and the MPA/JHU catalog. These differences are mainly due to the difference in the methods used to measure the continuum.\\footnote{Our code measures the continuum in the sidebands bracketing the emission line, while the MPA/JHU code measures it using a 200 pixel median smoothing of the emission-line-subtracted continuum.} However, the root cause of the systematic variation with redshift in both cases is due to a small-scale ($\\sim10$\\AA) systematic flux calibration residual, as we will demonstrate in this paper.   \\begin{figure*} \\begin{center} \\includegraphics[totalheight=0.35\\textheight]{n2s2_z_mine.ps} \\includegraphics[totalheight=0.35\\textheight]{n2s2_z_mpa.ps} \\caption{Left panels show the distribution of the \\niibw/\\siiw\\ flux ratio (upper panel) and $D_n(4000)$ (lower panel) for LINERs as a function of redshift. The sample is volume-limited with $M_{^{0.1}r} < -20.7$. LINERs are selected using the criteria described by \\cite{Kewley06}, except that no criterion involving \\oiw\\ is required. The dark points with tiny error bars show the median of the distribution. The error bar indicates the error of the median, computed using the formular given by \\cite{Beers90}. The wiggly systematic variations with redshift are unphysical. The right panels are similar to the left but using measurements from the MPA/JHU DR7 catalog. The exact measurement methods differ but they all show significant systematic variations with redshift, which are due to a small-scale systematic flux calibration residual. } \\label{fig:diagnosis_before} \\end{center} \\end{figure*}   In this paper, we will establish a new method for improving the flux calibration accuracy by an order of magnitude using coadded red galaxy spectra. Because red galaxies have very uniform spectra and evolve very slowly with redshift, when coadded they provide a high signal-to-noise standard. In addition, because this standard can be found at any redshift, no wavelength will be at disadvantage in the calibration due to features in the standard. We will show that this method can achieve a relative flux calibration accuracy of 0.1\\% on scales of a few hundred of Angstroms in SDSS.   In Section 2, we show how to measure the flux calibration residual in the SDSS using stacked red galaxy spectra, and test the improved spectrophotometry. We discuss potential improvements in Section 3 and summarize in Section 4. ",
        "conclusions": "We have demonstrated that the stacked spectrum of carefully selected red-sequence galaxies can be used as a spectrophotometric standard to produce an accurate relative flux calibration for large spectroscopic surveys. For SDSS, we have achieved the accuracy of $\\sim0.1\\%$. This method can be applied to surveys like the Baryon Oscillation Spectroscopic Survey (BOSS; \\citealt{Eisenstein11}), as long as the survey covers a sufficiently large number of red galaxies. In principle, other types of galaxies could also be utilized provided stacking a large enough sample can produce a reasonably stable spectrum over some wavelengths, which can be verified empirically, as done in Figure~\\ref{fig:ratio_restframe}.  However, this method is not meant to replace the traditional calibration methods using standard stars. To obtain best calibration, we still recommend the practice of observing standard stars simultaneously with science targets, as done in SDSS. They are essential for getting the smooth component in the spectrophotometry and for correcting for those extinction sources that vary on short timescales, such as the molecular absorption in the atmosphere. Our method is meant to be complementary and to improve a reasonably well calibrated spectrophotometry to the accuracy required for certain science applications, such as the study of weak line emission in red galaxies, Ly$\\alpha$ forest \\citep{Slosar11}, and any studies requiring high accuracy spectral index measurements."
    },
    "1107/1107.5106_arXiv.txt": {
        "abstract": "We analyze the three-year \\sdssii\\ Superernova (\\sn) Survey data and identify a sample of $1070$ photometric \\snia\\ candidates based on their multi-band light curve data.  This sample consists of \\sn\\ candidates with no spectroscopic confirmation, with a subset of 210 candidates having spectroscopic redshifts of their host galaxies measured, while the remaining 860 candidates are purely photometric in their identification.  We describe a method for estimating the efficiency and purity of photometric \\snia\\ classification when spectroscopic confirmation of only a limited sample is available, and demonstrate that \\snia\\ candidates from \\sdssii\\ can be identified photometrically with $\\sim 91\\%$ efficiency and with a contamination of $\\sim 6\\%$.  Although this is the largest uniform sample of \\sn\\ candidates to date for studying photometric identification, we find that a larger spectroscopic sample of contaminating sources is required to obtain a better characterization of the background events.  A Hubble diagram using \\sn\\ candidates with no spectroscopic confirmation, but with host galaxy spectroscopic redshifts, yields a distance modulus dispersion that is only $\\sim 20 - 40$\\% larger than that of the spectroscopically-confirmed \\snia\\ sample alone with no significant bias.  A Hubble diagram with purely photometric classification and redshift-distance measurements, however, exhibit biases that require further investigation for precision cosmology. ",
        "introduction": "\\label{section_intro}  Measurements of luminosity distances to nearby Type~Ia Supervova (\\snia) \\citep{phillips93, hamuy96a} and their distant counterparts have played a central role in modern cosmology and the remarkable discovery of an accelerating universe \\citep{riess98, perlmutter99}.  Many dedicated supernova (\\sn) surveys and follow-up programs have since then acquired light curves and spectra for several thousands of \\sn\\ in various redshift ranges: 1) at $z \\la 0.1$ by the Lick Observatory Supernova Search \\citep{filippenko01, ganeshalingam10}, the CfA monitoring campaign \\citep{riess99, jha06a, matheson08, hicken09}, SNFactory \\citep{bailey09}, Carnegie Supernova Project Low-$z$ Program \\citep{contreras09, folatelli10}, the Palomar Transient Factory \\citep{rau09, law09}, and the Panoramic Survey Telescope and Rapid Response System (\\panstarrs\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu/public}}); 2) the \\sdssii\\ \\sn\\ Survey in the intermediate redshift interval $0.1 \\la z \\la 0.3$ \\citep{frieman08, sako08}; 3) the highest-redshift range observable from the ground at $0.3 \\la z \\la 1$ by the Supernova Legacy Survey \\citep[\\snls;][]{astier06, guy10, conley11}, the \\essence\\ \\sn\\ Survey \\citep{miknaitis07, wood-vasey07}, the Carnegie Supernova Project High-$z$ Program \\citep{freedman09}; and finally 4) $z \\ga 1$ \\snia\\ from space using the Hubble Space Telescope \\citep{riess04a, riess07, dawson09}.  Many future surveys, such as the Dark Energy Survey \\citep[\\des;][]{flaugher10} and the Large Synoptic Survey Telescope \\citep[\\lsst;][]{lsst09}, with deeper and more wide-field imaging capabilities will probe much larger volumes of the universe allowing discoveries of thousands to tens of thousands of high-redshift \\sn\\ candidates each year.  Spectroscopic follow-up observations of these large, faint \\sn\\ samples will require prohibitively large time allocations with existing instruments.  Studies of \\sn\\ properties and cosmology will, therefore, necessitate a photometric determination of the \\sn\\ type, cosmological redshift, and the luminosity distance from light curves with possibly a limited subsample with spectroscopic confirmation and redshift measurements.  Various methods for photometrically classifying \\sn\\ have been discussed in the literature.  Optical and \\uv\\ colors near maximum light, for example, have been used to distinguish \\snia\\ from core-collapse \\sn\\ \\citep{pskovskii77, poznanski02, panagia03, riess04b, johnson06}. \\citet{poznanski07a} have developed a Bayesian method that classifies \\sn\\ using only a single epoch of photometry \\citep[see also,][]{kuznetsova07, rodney09}.  Template-fitting methods have been employed for spectroscopic targetting of active \\sn\\ candidates \\citep{sullivan06, sako08}.  \\citet{sullivan06} have performed an analysis to identify a sample of photometric \\snia\\ candidates from the first year of the Supernova Legacy Survey.  \\citet{dahlen04}, \\citet{poznanski07b}, \\citet{dahlen08}, \\citet{dilday08}, \\citet{dilday10}, \\citet{rodney10b}, and \\citet{graur11} have also used photometric classification to measure \\sn\\ rates as a function of redshift.  Although an efficient photometric \\sn\\ classifier is crucial for a successful spectroscopic follow-up program and also for understanding the bias in the spectroscopic sample, the ability to estimate both the efficiency and purity of the selected sample is also important for understanding, for example, possible biases in distance measurements and studies of \\sn\\ rates.  Clearly, the efficiency can be improved by compromising purity, and vice versa, and the requirements may vary depending on the type of study involved.  In addition to photometrically identifying \\snia\\ candidates, redshifts as well as luminosity distances can be inferred from the same multi-band light curve data.  These studies of {\\it \\sn\\ cosmology without spectroscopy} have been pioneered by \\cite{barris04} and carried out more recently by a number of authors.  \\citet{palanque10}, \\citet{kessler10a}, and \\citet{rodney10a} for example, study the quality of photometric redshifts on large samples of existing data.  \\citet{rodney10a} also construct a photometry-only Hubble diagram of the first-year \\sdssii\\ and \\snls\\ spectroscopically-confirmed \\snia\\ using their Supernova Ontology with Fuzzy Templates (\\soft) method. Others show comparisons of measured and input redshifts primarily from simulations \\citep{kim07, kunz07, wang07a, wang07b, gong09, scolnic09}.  The accuracy and precision of the measured parameters depend on many observational factors including the statistical quality of the observed light curves, surface brightness of the underlying host galaxy, photometric calibration, wavelength coverage, the number of filter bandpasses, and the observing cadence.  Other non-observational factors that might affect the measurements are the quality of the light curve models, assumptions on the dust properties and intrinsic \\sn\\ colors, as well as priors used in the fits.  The photometric redshift uncertainty on any individual \\sn\\ is obviously larger than a typical spectroscopic redshift error, but a substantially larger number of {\\emph{unbiased}} redshift and distance measurements made possible photometrically might be able to provide competitive constraints on cosmological parameters with future large-scale surveys.  Some of the existing softwares and algorithms, including the one presented in this paper, were recently used to participate in the Supernova Photometric Classification Challenge \\citep{kessler10b}, a public competition for classifying \\sn\\ light curves.  The authors of the challenge released a large number of simulated \\sn\\ light curves of undisclosed types and a small ``spectroscopic'' sample with known redshifts and types for training.  Participants of the challenge submitted their classifications as well as photometric redshifts if available.  The algorithm presented here achieved the highest overall figure of merit, though there is significant room for improvement.  This paper focuses on understanding these issues using an improved implementation of existing methods and through analysis of a much larger sample of \\sn\\ candidates for testing.  We use the three-year \\sdssii\\ \\sn\\ Survey data as our test bed to identify photometric \\snia\\ candidates with realistic estimates of sample purity.  The description of the photometric classification algorithm and the spectroscopic and photometric \\sn\\ samples from \\sdssii\\ are presented in \\S\\ref{sec:sncand}.  The procedures for estimating the \\snia\\ typing efficiency and purity using the spectroscopic sample are described \\S\\ref{sec:fom} and \\S\\ref{sec:eff}.  The properties of the photometric \\snia\\ candidates identified are described in \\S\\ref{sec:photo-ia}.  The quality of the light curve photometric redshifts is discussed in \\S\\ref{sec:redshift}.  Comparisons with simulations are shown in \\S\\ref{sec:sims}.  Finally, our results are summarized in \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have identified $1070$ photometric \\snia\\ candidates from the \\sdssii\\ \\sn\\ Survey data.  This sample is more than three times larger than the spectroscopically confirmed \\snia\\ sample with good light curves, and is estimated to include $\\sim 91$\\% of all \\snia\\ candidates detected by the survey with a purity of $\\sim 94$\\% ($\\sim 6$\\% contamination).  This estimate of the purity, however, is based on a limited number of spectroscopically confimred \\ccsn, most of which are nearby, bright events and are therefore not representative of the majority of the contaminating events.  As shown in Figure~\\ref{fig:snr}, the majority of our photometric candidates have peak $r$-band \\signoi $< 10$, where we have only a handful of spectroscopic \\sn\\ candidates.  To obtain a better characterization of the contaminating sources, confirmation is needed for a much larger sample of faint \\ccsn\\ that are comparable in apparent brightness to the photo-Ia sample.  As also advocated by \\citet{richards11}, future surveys that rely on photometric identification should obtain spectra of \\sn\\ candidates over the full range of the \\signoi\\ of the photometric candidates of interest.  The Hubble digram with photometric classification and host galaxy spectroscopic redshift priors show a slight increase in scatter over the confirmed \\snia\\ sample, which is consistent with them being due to mis-classified \\ccsn.  There is no significant redshift-dependent offset in the derived distances compared to the conf-Ia sample.  Simulations confirm these findings.  Photometric redshifts estimated from the multi-band light curves are unbiased below $z \\sim 0.2$ with an rms dispersion of $\\sigma_z \\sim 0.05$. There is a redshift-dependent bias above $z \\sim 0.2$ where the mean redshift difference $\\langle z_{\\mathrm{lc}} - z_{\\mathrm{photo}} \\rangle$ is between $-0.04$ and $-0.02$.  The rms dispersion is $\\sigma_z \\sim 0.05 - 0.10$.  The Hubble diagram of the photo-Ia sample also exhibits outliers and redshift-dependent biases.  Although the distance and redshift accuracies at present are not yet sufficient for cosmology, the large sample can still be used for studies of the \\snia\\ rate as a function of redshift, correlations between \\sn\\ light curves and host galaxy properties, and other studies that do not involve joint constaints on both redshift and distance.  We conclude that cosmology with future large-scale \\sn\\ surveys should at the minimum measure host galaxy spectroscopic redshifts for the Hubble digram.  A subset of the \\sn\\ candidates must be observed spectroscopically to study the photometric classification efficiency and purity.  Spectroscopy should target candidates with \\signoi\\ down to the magnitude limit where photometric classification is expected to work.  Cosmology with photometry alone, however, requires further investigation with realistic simulations in order to understand and characterize their systematic biases and uncertainties, and how they depend on the \\snia\\ candidate selection criteria."
    },
    "1107/1107.2510_arXiv.txt": {
        "abstract": "{A large sample of T Tauri stars exhibits optical jets, approximately half of which rotate slowly,   only at ten per cent of their breakup velocity. The disk-locking mechanism has been shown to be inefficient to explain this observational fact.} {We show that low mass accreting T Tauri stars may have a strong stellar jet component that can effectively brake the star to the observed rotation speed.} {By means of a nonlinear separation of the variables in the full set of the MHD equations we construct semi-analytical solutions describing the dynamics and topology of the stellar component of the jet that emerges from the corona of the star.} {We analyze two typical solutions with the same mass loss rate but different magnetic lever arms and jet radii. The first solution with a long lever arm and a wide jet radius effectively brakes the star and can be applied to the visible jets of T Tauri stars, such as RY Tau. The second solution with a shorter lever arm and a very narrow jet radius may explain why similar stars, either Weak line T Tauri Stars (WTTS) or Classical T Tauri Stars (CTTS) do not all have visible jets. For instance, RY Tau itself seems to have different phases that probably depend on the activity of the star. } {First, stellar jets seem to be able to brake pre-main sequence stars with a low mass accreting rate. Second, jets may be visible only part time owing to changes in their boundary conditions.  We also suggest a possible scenario for explaining the dichotomy between CTTS and WTTS, which rotate faster and do not have visible jets. } ",
        "introduction": "\\subsection{The angular momentum problem, observational facts}  Observations of star-forming regions show that several T Tauri stars are associated with well collimated jets. These jets are made of strongly accelerated plasma that tra\\-vels at a few hundred km/s \\citep{Bally09} and is magnetically collimated at large distances \\citep{Dougadosetal00}. The jets are usually associated with Classical T Tauri Stars (CTTS), which are low mass stars $\\le 2 M_{\\odot}$ in their late stages of pre-main sequence evolution, showing strong evidence of the presence of a surrounding accretion disk. The outflow is usually thought to be ejected from the Keplerian disk (e.g. \\citeauthor{CabritAndre91}, \\citeyear{CabritAndre91}). However, a significant part of the jet, if not all, may be ejected by the star itself, at least for the lower mass accreting T Tauri stars.  Observations also reveal that approximately half of the T Tauri stars rotate slowly, at about $10\\%$ or less of their breakup speed \\citep{Mattetal10}. This indicates that a very efficient mechanism is at work to remove the angular momentum in these stars; the nature of this mechanism is still controversial.  \\cite{Kundurthyetal06},  see also \\cite{Edwardsetal93},  claimed that  CTTS are slow rotators with circumstellar disks, while WTTS are fast rotators without accretion disks. However, the association of CTTS with slow rotators and WTTS with fast ones is not as yet observationally confirmed. Moreover, \\citeauthor{Stassunetal99} (\\citeyear{Stassunetal99}, \\citeyear{Stassunetal01}) and \\cite{Rebulletal04} could not find a strong correlation of fast rotators with stars that have dispersed their disks. Conversely, new results  (\\citeauthor{Rebulletal06}, \\citeyear{Rebulletal06}; \\citeauthor{Herbstetal07}, \\citeyear{Herbstetal07}; \\citeauthor{CiezaBaliber07}, \\citeyear{CiezaBaliber07}) found a bimodal distribution where stars with disks rotate slower than stars without disks. However, the two distributions clearly overlap. The prevailing theory has been that magnetic interaction between the star and its circumstellar disk regulates the stellar rotation periods. Attempts to observationally confirm  this conjecture have produced mixed results. Hence, from the observational point of view, the presence of an accretion disk may not be enough to explain stellar angular momentum removal.  In order to understand slow stellar rotators, \\cite{Schatzman62} suggested that the magnetic braking of a stellar wind would be sufficient even with a low mass loss rate, and this has been explored by several authors (e.g. \\citeauthor{WeberDavis67}, \\citeyear{WeberDavis67}, \\citeauthor{Mestel68a}, \\citeyear{Mestel68a}, \\citeyear {Mestel68b}). Following this track, it is essential to address stellar wind models in detail, with and without the presence of an accretion disk, to see how efficient the wind is in extracting angular momentum from the star. Moreover,  recent observations confirm  that stellar winds are present in at least $60\\%$ of CTTS \\citep{Kwanetal07}, as was already suggested by various authors (e.g. \\citeauthor{Edwardsetal03}, \\citeyear{Edwardsetal03}).  \\cite{Decampli81} showed that  {\\it thermal} stellar winds cannot support jets with a high mass loss rate, simply because this would require too high temperatures to be physical. For low mass loss rates, less  than $10^{-9}M_\\odot/yr$, he showed that a thermally driven or line-driven radiative wind can support the formation of the jet. For a higher mass loss rate, such a driving is still possible provided that there is an extra supply of pressure from Alfv\\'en waves. If the temperature of the wind is around one million degrees, the pressure can drive the jet up to 300  km/s. However, the pressure may exceed the thermal pressure by a large amount  if it  includes ram pressure or turbulent magnetic pressure.  The situation is similar to the well known case of the solar wind, where to obtain wind speeds of several hundred km/s with a temperature of a few million degrees one requires extra pressure from either turbulence or kinetic effects. These effects yield an effective  temperature of around ten million degrees (see \\citeauthor{Aibeoetal07}, \\citeyear{Aibeoetal07} and references therein).  Ignoring the important question regarding the nature of the pressure and following the early study of \\cite{Decampli81}, most authors disregarded stellar wind models up to now. This explains why disk wind models have been extensively studied.  \\subsection{Disk wind theory}  In disk wind models, the main source of acceleration is the magneto-centrifugal driving, as proposed originally by \\cite{BlandfordPayne82}, who presented the first MHD radial self-similar solutions extending the earlier hydrodynamical solutions of \\citeauthor{BardeenBerger78} (\\citeyear{BardeenBerger78}) for galactic winds. Since then, their original model has been improved in various ways (e.g., \\citeauthor{Lietal92}, \\citeyear{Lietal92}; \\citeauthor{Contopoulos94}, \\citeyear{Contopoulos94} ; \\citeauthor{VlahakisTsinganos98}, \\citeyear{VlahakisTsinganos98}; \\citeauthor{Vlahakisetal00}, \\citeyear {Vlahakisetal00}). In particular \\citeauthor{Ferreira97} (\\citeyear{Ferreira97}) has shown that a consistent treatment of the connection with the disk is possible.  From the theoretical point of view however, self-similar disk wind models are not consistently describing the inner part of the jet close to its axis \\citep{Graciaetal06}. The solution has to be cut at the inner and outer edges of the disk. From the observational point of view, cold disk winds turned out to be too rapid and too light (\\citeauthor{Garciaetal01b}, \\citeyear {Garciaetal01b}). This last point can be overcome with warm disk solutions and explains the high speed, dense powerful jets of early TTauri stars like DG Tau, as well as their rotation (\\citeauthor{Ferreiraetal06}, \\citeyear{Ferreiraetal06}). DG Tau is probably one of the most powerful sources with an accretion rate of $10^{-6}M_\\odot/$yr and a corresponding mass loss rate that is about 10\\% of this value \\citep{Hartiganetal95}. In this case, disk wind solutions seem more adapted. Recent observations suggest that DG Tau is variable on short timescales with an accretion rate dropping down to $10^{-7}M_\\odot/$yr \\citep{Becketal10}.  Simultaneously, several numerical simulations have investigated the jet launching of these disk winds (e.g., \\citeauthor{OuyedPudritz97}, \\citeyear{OuyedPudritz97}; \\citeauthor{Ustyugovaetal00}, \\citeyear{Ustyugovaetal00}; \\citeauthor{Krasnopolskyetal03}, \\citeyear{Krasnopolskyetal03}; \\citeauthor{CasseKeppens04}, \\citeyear{CasseKeppens04}). By investigating various boundary conditions and initial configurations, they confirmed that disk winds could indeed be accelerated and collimated by their own magnetic field. More recently several authors (\\citeauthor{Matsakosetal08}, \\citeyear{Matsakosetal08}; \\citeauthor{Stuteetal08}, \\citeyear{Stuteetal08}; \\citeauthor{Matsakosetal09}, \\citeyear{Matsakosetal09}) have shown that analytical solutions such as the radially self-similar ones are structurally stable. This result indicates that these solutions are good complementary tools to numerical simulations.  According to theoretical arguments \\citep{Pudritz&Norman86} and numerical simulations (e.g.  \\citeauthor {Melianietal06}, \\citeyear{Melianietal06}), disk winds can remove most of the angular momentum of the accreting plasma. However, disk winds need the presence of an inner stellar jet component to be globally consistent.  From the theoretical point of view,  it was initially proposed that fieldlines of the stellar magnetosphere anchored into the disk could efficiently brake the central star, in the so-called disk-locking mechanism proposed by \\cite{ChoiHerbst96}. In principle a fieldline rooted into the star will slow down (respectively speed up) the central star rotation, if it connects to a region of the disk rotating at lower rotational speed (respectively higher speed) after (respectively before) the corotation radius. Using a simple analytical model based on the Ghosh \\& Lamb mechanism (\\citeyear{GhoshLamb79a}, \\citeyear {GhoshLamb79b}), \\citeauthor{MattPudritz05} (\\citeyear{MattPudritz05}) showed however that the disk-locking efficiency is reduced because it leads to an opening of the fieldlines that disconnects the disk from the star. Instead they proposed that angular momentum can be lost via accretion-powered stellar winds provided the mass loss rate is 1 to 10 \\% of the accretion rate as usually inferred from observations \\citep{Calvet98}. Their results were confirmed via numerical simulations (\\citeauthor{MattPudritz08a}, \\citeyear{MattPudritz08a}, \\citeyear{MattPudritz08b}; \\citeauthor{Mattetal10}, \\citeyear{Mattetal10}). \\cite{Kuekeretal03}, using similar simulations, concluded that even when the disk-magnetosphere connection is not disrupted, the magnetic field could slow down the central object but the dominant disk torque still spins up the star. As mentionned by \\cite{Romanovaetal09}, these conclusions are not definitive because long-term simulations are required to establish the validity of this result. Ultimately the authors insist on the essential role of the axial jet. \\cite{ZanniFerreira09} performed various additional simulations where the disk-locking is not disrupted by the building of large toroidal field. However, even in this situation, the magnetic accretion torque is not sufficient to cause an efficient spin-down of the star.  \\subsection{Stellar jet  theory}  As a conclusion of the precedent subsection, the presence of a pure stellar wind in addition to a possible outer disk wind is required to effectively remove the excess of  stellar angular momentum.  Besides, more evolved T Tauri stars like RY Tau seem to have weaker jets with accretion rates on the order of only a few $10^{-8}M_\\odot/$yr, which may be an indication that disk winds are not always absolutely necessary for the weaker and more evolved jets. We see that with a mass loss rate of only 10\\% of this value, these jets are more difficult to analyze. From spectroscopic observations of RY Tau, \\citeauthor{GomezVerdugo01} (\\citeyear {GomezVerdugo01}, \\citeyear{GomezVerdugo07}) suspected the presence of a small-scale pure stellar jet because the UV emission lines originate in a region that is too small to be produced by a disk wind. New observations (\\citeauthor{StongeBastien08}, \\citeyear{StongeBastien08}; \\citeauthor{AgraAmboageetal09}, \\citeyear{AgraAmboageetal09}) clearly show a microjet in this faint object. It is interesting to investigate if these (micro-)jets can be modeled through stellar winds or a combination of a stellar wind and a sub Keplerian disk wind.  This combination of a two-component outflow is under numerical investigation (\\citeauthor{Mattetal03}, \\citeyear{Mattetal03}; \\citeauthor{Koide03}, \\citeyear{Koide03}; \\citeauthor{Kuekeretal03}, \\citeyear{Kuekeretal03}; \\citeauthor {Matsakosetal08}, \\citeyear{Matsakosetal08}). Recent simulations (\\citeauthor{Melianietal06}, \\citeyear{Melianietal06}; \\citeauthor{Romanovaetal09}, \\citeyear{Romanovaetal09}; \\citeauthor{Matsakosetal09}, \\citeyear {Matsakosetal09}; \\citeauthor{Fendt09}, \\citeyear{Fendt09}) also show the interplay of mixing between the two components. Several authors, e.g. \\cite{Kuekeretal03}, \\cite{Romanovaetal09} and \\cite{Fendt09}, have shown in their numerical simulations that the disk-magnetosphere connection produces strong intermittent outflows, the results of flares and reconnection. This is somewhat similar to coronal mass ejections in the solar wind and in microquasars. Though important to model, those outflows are not necessarily more important on the long-term duration compared to the continuous underlying steady flow, as suggested by \\cite{Romanovaetal09}. Moreover, \\cite{Matsakosetal09} have shown that the variability of the source does not destroy the continuous steady jet. Indeed, they use the first solution presented in this paper, as initial condition for the inner spine jet of their simulations, but use a polytropic equation of state. In their previous paper \\citep{Matsakosetal08} they have shown that changing from non polytropic to polytropic would reduce the size of the radius of the jet, without drastically affecting the outflow behavior, however.  However, numerical simulations are usually time-consuming to such an extent that they cannot explore a wide range of parameters. Moreover, they do not really start ejecting mass directly from the star but rather at a given height between the sonic and the Alfv\\'{e}n surfaces. Therefore, it is of absolute interest to model the stellar jet that originates in the central star.  All these problems can be overcome by studying MHD solutions obtained via a nonlinear separation of the variables. These so-called self-similar solutions contain as special cases all known MHD outflow solutions (e.g., \\citeauthor{VlahakisTsinganos98}, \\citeauthor{STT02}, \\citeyear{STT02}, hereafter STT02). As shown in \\cite{ST94}, hereafter ST94, these self-similar models can be seen as a combination of stellar wind and X-wind models. Conversely to formal X-winds (\\citeauthor{Shuetal94}, \\citeyear{Shuetal94}; \\citeauthor{Shangetal02}, \\citeyear{Shangetal02}), they do not require that all fieldlines emerge from a single point of the disk in the form of a fan. Instead there is a smooth transition between the stellar wind and the disk jet. They are also shown to be structurally stable (\\citeauthor{Matsakosetal08}, \\citeyear{Matsakosetal08}; \\citeyear{Matsakosetal09}). This was far from obvious because the heating function is not polytropic as usually in most analytical solutions. In a series of papers -- \\cite{STT99} hereafter STT99, STT02, and \\cite{STT04}, hereafter STT04 -- we have systematically explored the full parameter space of this problem. In ST94 we have shown that the double disk and star component was promising. However, the  solutions presented there had mass loss rates that were too low. Here we present solutions where the stellar jet is both under-dense and under-pressured compared to the disk wind in the launching region. These solutions have the advantage to adapt themselves not only to CTTS, which are connected with the surrounding accretion disk, but also to WTTS where there is { either no accretion disk or no connection with the disk itself} \\citep{Walteretl88, Bertout89}. Although winds from WTTS are difficult to measure, mass loss rates between $10^{-11}$ and  $10^{-10}M_\\odot/$yr have been obtained by  \\cite{Andreetal92}. Winds of WTTS might be similar to the solar wind in terms of collimation (see \\citeauthor{Aibeoetal07}, \\citeyear{Aibeoetal07}, and references therein). However, these young stars have higher mass loss rates and rotational speeds than our  sun and might as well have strongly collimated jets { independently of the presence of an accretion disk}. Their rotation periods range from  0.6  to 24 days with peaks near two and eight days \\citep{Herbstetal07} while the solar rotational period is 24.5 day. Thus,  WTTS may well produce self-collimated jets that nevertheless are so weak that they cannot be detected.  We will use the ST94 meridionally self-similar model to examine the contribution of the stellar component to the overall jet and how efficiently it brakes the star. We first explore how  the model parameters can be constrained from the observations of CTTS jets (Sect. \\ref{sec.3}). Then, we focus our attention on two different classes of solutions (Sections \\ref{sec.4} and \\ref{sec.5}) and explore the efficiency of these winds to spin down the central object. In the final section (Sect. \\ref{sec.6}) we discuss how these two classes of solutions may explain the dichotomy between CTTS and WTTS and how they can be adapted to also model multiphases of T Tauri faint jets, such as in RY Tau. ",
        "conclusions": "\\label{sec.6} We have used a simple semi-analytical MHD model (ST94) to show that the low rotational velocity of T Tauri stars can be understood by considering the effect of stellar jets to remove their angular momentum. In particular, we modeled stellar jets from CTTS with a mass loss rate consistent with observations. The solutions for collimated magnetized outflows are obtained via a nonlinear separation of the variables in the governing full set of the MHD equations. Two main results are obtained.   {\\it First}, jets can efficiently brake the central star and remove the bulk of its angular momentum over a timescale of a million years, at least for low mass accreting stars. This corresponds to the typical lifetime of  a star in the CTTS phase. Furthermore, with this mechanism the star can even slow down within 0.6 million years, if the disk does not transport angular momentum onto the star. This fairly short slow-down time shows the efficiency with which a magnetized outflow can remove angular momentum from a rapidly rotating young stellar object. This result strongly suggests that stellar jets may explain the low rotational speeds observed in those objects, independently of any disk locking. This mechanism could be complementary to other angular momentum removal mechanisms occurring at the X- wind emerging from the interaction zone between the disk and the magnetosphere. Our solution fairly well reproduces the mass loss rate, terminal speed and rotation rate of the jet. The effective plasma temperature in the jet is reasonable provided that non thermal processes, such as Alfv\\'en waves, are at work in the plasma of the jets. These physical processes are likely to take place in the jets if one considers that the magnetic activity of CTTS is comparable to, or higher than, that of our own Sun.   {\\it Second}, we also analyzed two interesting MHD solutions. The first corresponds to a  gradually cylindrically collimated wide wind without oscillations in its width. The second is a narrower outflow that refocuses toward the jet axis and also oscillates in width. Note that these two very different solutions were obtained by a slight change of the parameters. It is a known result in nonlinear equations (e.g. Parker solutions, or, their MHD generalization). Correspondingly, we discussed the applicability of these two solutions in the context of the RY Tau jet, of which recent observations indicate that the plasma outflow from the rotating and magnetized star may have an intermittency exhibiting at least two phases. During one phase the cylindrical wind is relatively wide and nonoscillating. During the other phase the outflow is narrow and oscillates. Oscillations may produce radiative shocks seen in various emission lines, such as the observed UV lines of RY Tau. \\cite{Matsakosetal09} have shown that the self-similar solutions used in this paper will not be destroyed by the shock formation or the disk-wind interaction.  From a more general perspective, this dichotomic behavior can be seen in the two classes, CTTS and WTTS, which may both actually have jets, but these jets are visible only CTTS. The connection of CTTS with their disks through their magnetospheres may increase the width of their jets such as to make them visible. The mass loss rate measured in CTTS is thought to decrease as the star evolves toward the stage of WTTS. In this later stage, direct observations are more difficult and upper limits around $10^{-10}M_\\odot/$yr can only be inferred. Weak T Tauri stars seem to have lost the gas component of their accretion disk or at least gas falling onto the star. Our second MHD solution suggests that even with the same mass loss rate, the jets of WTTS may be invisible with the present angular resolution because of their narrowness. The lack of accretion disk prevents the formation of disk winds and the possibility of higher mass loss rates. The lower jet width may be caused either by a lower stellar magnetic activity (see \\citeauthor{Vidottoetal09}, \\citeyear{Vidottoetal09}, \\citeyear{Vidottoetal10}) or by a lack of magnetic connection into the disk. The magnetic braking time in this case would be on the order of 10 million years, which is indeed the time a T Tauri star is expected to spend in the weak-line stage.  Our modeling suggests that CTTS may produce visible jets in their active phase that would efficiently brake the star. Indeed, in the nonoscillating solution the magnetic field dominates the dynamics in the jet. Near the star, the magnetic field has a dipolar structure and the wind is more efficient in this case in extracting stellar angular momentum. On the other hand, WTTS may have jets as well, which are invisible because they are too narrow. These jets would be weaker because of the lack of direct connection with the accretion disk through a large dead zone."
    },
    "1107/1107.1254_arXiv.txt": {
        "abstract": "We examine predictions for the quasar luminosity functions (QLF) and quasar clustering at high redshift ($z \\ge 4.75$) using {\\it MassiveBlack}, our new hydrodynamic cosmological simulation which includes a self-consistent model for black hole growth and feedback.  We show that the model reproduces the Sloan QLF within observational constraints at $z \\ge 5$.  We find that the high-z QLF is consistent with a redshift-independent occupation distribution of BHs among dark matter halos (which we provide) such that the evolution of the QLF follows that of the halo mass function.  The sole exception is the bright-end at $z=6$ and 7, where BHs in high-mass halos tend to be unusually bright due to extended periods of Eddington growth caused by high density cold flows into the halo center.  We further use these luminosity functions to make predictions for the number density of quasars in upcoming surveys, predicting there should be $\\sim 119 \\pm 28$ ($\\sim 87 \\pm 28$) quasars detectable in the F125W band of the WIDE (DEEP) fields of the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey (CANDELS) from $z=5-6$, $\\sim 19 \\pm 7$ ($\\sim 18 \\pm 9$) from $z=6-7$, and $\\sim 1.7 \\pm 1.5$ ($\\sim 1.5 \\pm 1.5$) from $z=7-8$. We also investigate quasar clustering, finding that the correlation length is fully consistent with current constraints for Sloan quasars ($r_0 \\sim 17 \\: h^{-1} \\: \\rm{Mpc}$ at $z=4$ for quasars above $m_i = 20.2$), and grows slowly with redshift up to $z=6$ ($r_0 \\sim 22 \\: h^{-1} \\: \\rm{Mpc}$).  Finally, we note that the quasar clustering strength depends weakly on luminosity for low $L_{\\rm{BH}}$, but gets stronger at higher $L_{\\rm{BH}}$ as the BHs are found in higher mass halos. ",
        "introduction": "Supermassive black holes are believed to be present at the center of most galaxies \\citep{KormendyRichstone1995} and are found to be correlated with the properties of their host galaxies \\citep{Magorrian1998, FerrareseMerritt2000, Gebhardt2000, Tremaine2002, GrahamDriver2007}.  These correlations provide strong evidence for a link between the growth of black holes and the evolution of their host galaxies, generally attributed to some form of quasar feedback \\citep{BurkertSilk2001, Granato2004, Sazonov2004, Springel2005, Churazov2005, KawataGibson2005, DiMatteo2005, Bower2006, Begelman2006, Croton2006, Malbon2007, CiottiOstriker2007, Sijacki2007, Hopkins2007}.   Perhaps the most fundamental statistical quantity in the study of quasars is the number density, often characterized as a Quasar Luminosity Function (QLF).  The QLF has been studied observationally \\citep{LaFranca2002, Fiore2003, Ueda2003, Barger2003, Croom2004, LaFranca2005, Cirasuolo2005, Richards2006, Brown2006, Silverman2008, Ebrero2009, Yencho2009}, as well as through simulations, with black holes modeled both semi-analytically \\citep{KauffmannHaehnelt2000, Volonteri2003, WyitheLoeb2003, Granato2004, Marulli2008, Bonoli2009} and directly incorporated in the simulation \\citep{Hopkins2005, HopkinsQLF2006, Hopkins2007, Marulli2009, DeGraf2010}.  Another fundamental quantity for quasar populations is the strength of quasar clustering, and its evolution with redshift.  Numerous observational studies have been done \\citep{LaFranca1998, Porciani2004, Croom2005, Shen2007, Myers2007, daAngela2008, Shen2009, Ross2009}, generally finding evidence for increasing clustering amplitude with redshift, in agreement with findings from simulations \\citep{Bonoli2009, Croton2009, DeGraf2010}.  Clustering is especially significant because one can use it to estimate the typical dark matter halos in which quasars are found simply by comparing the clustering strength of quasars to that of dark matter halos. In particular, the luminosity dependence of the clustering (if any) can help determine if the bright and faint quasars populate the same halos, suggesting they may be the same population of quasars at different phases of their lives \\citep[see, e.g.][]{Hopkins2005, Hopkins2005b, Hopkins2005c, Hopkins2005d, Hopkins2006b} or different halo masses, suggesting they are fundamentally different quasar populations \\citep[see, e.g.][]{Lidz2006}.  Thus both QLF and quasar clustering can be used to investigate the populations of quasars being observed, how they populate dark matter halos, and how the typical luminosities correlate with hosts.  However, investigations at high redshift have been extremely difficult.  Simulations have proved difficult due to the volumes needed to produce high redshift quasars in sufficient numbers, limiting the statistical investigations.  Similarly, the difficulty observing such distant objects limits the current number of observed high redshift quasars such that the $z \\sim 6$ faint-end QLF is determined by only a few objects \\citep{WillottCFHQS2010}, and high-redshift quasar clustering is limited to lower redshift \\citep[e.g. $z \\sim 4$ by][]{Shen2009}.  However, upcoming surveys such as the James Webb Space Telescope (JWST) and the Cosmic Assembly Near-infrared Deep Extragalactic Legacy Survey (CANDELS) have the potential to drastically improve the high-z observations, making now an ideal time to make predictions as to the statistics of high redshift quasars.  In this paper we use a new large-scale hydrodynamic cosmological simulation % to study the earliest supermassive black holes, focusing on their luminosity and clustering properties.  We take advantage of this simulation's large volume (it is the largest cosmological simulation which directly models the growth and evolution of black holes) to investigate the statistical properties of the earliest supermassive black holes, focusing on the luminosity function and correlation length.  We use these quantities to compare with current observational measurements, to make predictions for upcoming surveys, and to investigate the implications of potential features which may be detected.  In Section \\ref{sec:Method} we describe the numerical sumulation used for our analysis.  In Section \\ref{sec:Results} we investigate the Quasar Luminosity Function (Section \\ref{sec:QLF}) and quasar clustering (Section \\ref{sec:Clustering}) and compare with observations, and our results are summarized in Section \\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  In this paper we have investigated high redshift black holes found within a new large-scale hydrodynamic simulation, focusing on the luminosity function and clustering behavior.  We model the QLF for $z \\ge 5$, probing black holes to luminosities up to $\\sim 10^{14} L_\\odot$.  We find reasonable agreement with observational data at $z=5$ and 6, generally falling within the variation between surveys, confirming that our model is capable of producing Sloan-type quasars at very early times as well as matching their observed abundances.  Using a HOD fit for the central AGN at $z=5$ (which we have provided), we find that the evolution in the QLF is well described by the evolution in the halo mass function (at least for the redshifts modeled by our simulation) without significant evolution in the occupation distribution, except for the high-luminosity end at $z=6-7$ where we find a significantly flatter luminosity function.  We postulate this flattening of the luminosity function to be a result of the largest black holes tending to undergo unusually rapid growth (and thereby producing unusually high luminosities) during these redshifts, a conclusion supported both by the difference in the HOD fits at these redshifts, and by direct investigation into the lightcurves of these black holes \\citep{DiMatteo2011}.  At lower redshift (by $z \\sim 5$), self-regulation suppresses the most massive BH growth, resulting in fainter bright-end BHs, and a corresponding steepening of the QLF.  In particular, we note that this increase in luminosity at $z \\sim 6-7$ should make observations of such high-z, luminous ($>10^{12} L_\\odot$) BHs easier than would otherwise be expected.  We used our luminosity function to provide estimates on the number density of detectable high redshift quasars for several upcoming surveys.  In particular, we expect the CANDELS survey to find quasars up to $z \\sim 7$ in both the WIDE and DEEP programs, with each probing a slightly  different region of the QLF.  At $z \\sim 6$ these programs should each detect dozens of quasars across a wide range of luminosities, drastically improving the observational constraints on the faint end of the high redshift QLF, particularly the faint-end slope.  In addition, we have provided our estimated number density at each redshift for several additional upcoming surveys, thereby providing the approximate number density of quasars to be found as well as the survey areas necessary to reach the highest redshifts.  We also investigate the clustering behavior of these high redshift quasars, finding luminosity dependent correlation lengths on the order of  $\\sim 10-14 \\: h^{-1} \\rm{Mpc}$.  Using our HOD model and the theoretical clustering of dark matter halos, we show the correlation length as a function of BH luminosity, finding relatively weak luminosity dependence at low $L_{\\rm{BH}}$, but with increasing $L_{\\rm{BH}}$-dependence at higher luminosities (where black holes are typically found in halos in the steep end of the halo mass function).  We also compare our predicted $r_0$ to high-redshift observations and find excellent agreement, with our simulation predicting $r_0 \\sim 15-20 \\: h^{-1} \\: \\rm{Mpc}$ for Sloan-type quasars.  We also roughly match the evolution of $r_0$ with redshift using a redshift-independent HOD, suggesting that high-redshift quasar clustering evolution is fully explained solely by the evolution in clustering of dark matter halos without a change in the typical host halo mass.  Finally, we note the effect a change in the BH occupation distribution from $z=5$ to 6 has on $r_0$, finding that it should suppress the clustering strength of high luminosity quasars at $z=6$.  Our limited sample size and evolution in the halo bias factor found in our simulation make it difficult to quantify the magnitude of this suppression, but we nonetheless conclude that some suppression should occur (though likely below the sensitivity of upcoming surveys)."
    },
    "1107/1107.1581_arXiv.txt": {
        "abstract": "{ The surface tension of cold and dense QCD phase transitions has appeared recently as a key ingredient in different astrophysical scenarios, ranging from core-colapse supernovae explosions to compact star structure. If the surface tension is low enough, observable consequences are possible. Its value is however not known from first-principle methods in QCD, calling for effective approaches. Working within  the framework of homogeneous nucleation by Langer, we discuss the steps that are needed to obtain the nucleation parameters from a given effective potential. As a model for deriving the effective potential for the chiral transition, we adopt the linear sigma model with constituent quarks at very low temperatures, which provides an effective description for the thermodynamics of the strong interaction in cold and dense matter, and predict a surface tension of $\\Sigma \\sim 5$--$15~$MeV/fm$^{2}$, well below previous estimates. Including temperature effects and vacuum logarithmic corrections, we find a clear competition between these features in characterizing the dynamics of the chiral phase conversion.}  \\FullConference{The many faces of QCD\\\\ November 1-5, 2010\\\\ Gent, Belgium}  \\begin{document}  One of the many faces of QCD that has been attracting continuous attention in the last years is how the matter dictated by the fundamental theory of strong interactions behaves at extreme conditions of temperature and/or density. The fact that the energy scales for the QCD phase transitions are currently under reach of different heavy-ion colliders creates an intense demand of theoretical and phenomenological development.  Over the past decade RHIC has established the formation of a new state of matter, characterized by a clear fluid behavior at the partonic level, differing from the predicted picture of the na\\\"ive weakly-interacting quark gluon plasma \\cite{Adams:2005dq}. LHC stretches the energy frontier also for heavy-ion collisions, bringing up the statistics and allowing for a careful study of fluctuations that might reveal the features of the initial state of the collisions.  The richest structure in the high energy QCD phase diagram seems however to be deep within the region with an asymmetry between the particle and antiparticle content of the medium. In this direction many experimental efforts are dedicated to create matter with a finite baryon-chemical potential: the aim is the search for the phase boundaries in the QCD phase diagram and its properties. The accelerators however cannot reach the domain of very low temperature and high chemical potential, which is also the region where theorists have less guidance for what QCD predicts, since first-principle lattice calculations are not possible \\cite{Laermann:2003cv}.  Astrophysical phenomena might however shed some light in this partially obscure region of the QCD phase diagram. Extremely high densities and relatively low temperatures are believed to exist in the core of ultracompact objects, so that it is probable that the QCD phase transitions play an important role in the structure and dynamics of such systems. Even though in general observations are much less controlable than laboratory experiments, the astronomical measurement techniques and instrumentation have much developed in the last decade and this might be the dawn of a precision era in which astrophysics really constrains nonperturbative QCD phenomena. One concrete example of astronomical observation that has reached the capability of constraining models for strong interactions in the nonperturbative regime is that of Ref. \\cite{nature} in which the error bars of less than $5\\%$ in the $2$ solar masses measurement rule out different equations of state.  The surface tension between partonic and hadronic phases of QCD seems to be a key quantity in connecting microscopic modeling of cold and dense QCD with astrophysical applications: recently it has been put forward that a low enough QCD surface tension should allow for possible observable consequences in different phenomena related to compact objects. It was shown for instance that deconfinement can happen during the early post-bounce accretion stage of a core collapse supernova event, which could result not only in a delayed explosion but also in a neutrino burst that could provide a signal of the presence of quark matter in compact stars \\cite{Sagert:2008ka}. However, as was discussed in detail in Ref. \\cite{Mintz:2009ay} (see also Ref. \\cite{Bombaci:2009jt} )  those possibilities depend on the actual dynamics of phase conversion, more specifically on the time scales that emerge. In a first-order phase transition, as is expected to be the case in QCD at very low temperatures, the process is guided by bubble nucleation (usually slow) or spinodal decomposition (``explosive'' due to the vanishing barrier), depending on how fast the system reaches the spinodal instability as compared to the nucleation rate \\cite{reviews}. Nucleation in relatively high-density, cold strongly interacting matter, with chemical potential of the order of the temperature, can also play an important role in the scenario proposed in Ref. \\cite{Boeckel:2009ej}, where a second (little) inflation at the time of the primordial quark-hadron transition could account for the dilution of an initially high ratio of baryon to photon numbers. Moreover, significantly different compact star structures are obtained if one considers the possibility of a layer of a quark-hadron mixed phase in the core of compact stars \\cite{Kurkela:2010yk}. A key ingredient in all these scenarios is, of course, the surface tension, since it represents the price in energy one has to pay for the mere existence of an interface between quark and hadron phases. If the cost in energy is too high, quark-hadron mixed phases will not be favourable and similarly nucleation time scales will be too long for these astrophysical phenomena to take place.  Computing the surface tension for the QCD transitions with first-principle methods is a prohibitively complicated task. It is a genuine nonperturbative quantity, in the sense that it cannot be obtained by a na\\\"ive expansion around one minimum of the effective action, and is also inaccessible by current first-principle lattice QCD simulations, due to the sign problem. One has therefore to resort to effective models. Estimates for the surface tension between a quark phase and hadron matter were considered previously in different contexts. In a study of the minimal interface between a color-flavor locked phase and nuclear matter in a first order transition, the authors of Ref. \\cite{Alford:2001zr} use dimensional analysis and obtain $\\Sigma\\sim 300~$MeV/fm$^{2}$ assuming that the transition occurs within a fermi in thickness. Taking into account the effects from charge screening and structured mixed phases, the authors of Ref. \\cite{Voskresensky:2002hu} provide estimates in the range of $50-150~$MeV/fm$^{2}$ but do not exclude smaller or larger values. These values of the surface tension are too high, rendering the nucleation of quark matter bubbles improbable in the context of compact objects.   We have estimated  \\cite{Palhares:2010be} the surface tension for the cold and dense chiral phase transition within a well-established chiral model (namely, the linear sigma model with constituent quarks (LSMq)) and assuming the homogeneous nucleation scenario \\cite{Langer:1967ax}. Our results indicate that the surface tension in the cold and dense regime can be considerably smaller than previous estimates, possibly allowing for the interesting observable astrophysical phenomena described above.  This paper is organized as follows. In Section \\ref{homnucl}, we delineate how, within the framework of homogeneous nucleation, it is possible to derive the nucleation parameters starting from an effective potential for the chiral order parameter $V_{\\rm eff}(\\sigma)$. Section \\ref{EffM} contains a short presentation of the effective model we adopted for the QCD chiral phase transition and presents the results for the effective potential $V_{\\rm eff}(\\sigma)$ for three different cases. Our results for the surface tension are discussed in Section \\ref{Res} and Section \\ref{Conc} contains final remarks. ",
        "introduction": " ",
        "conclusions": "\\label{Conc}  We have quantified approximately the homogeneous nucleation in the linear sigma model with quarks in the $\\overline{\\rm MS}$ scheme at zero and low temperature and finite quark chemical potential, including vacuum and medium fluctuations. All the relevant quantities were computed \\cite{Palhares:2010be} as functions of the chemical potential, the key function being the surface tension.  The inclusion of temperature effects and vacuum logarithmic corrections revealed a clear competition between these features in characterizing the dynamics of the chiral phase conversion, so that if the temperature is low enough the consistent inclusion of vacuum corrections could help preventing the nucleation of quark matter during the collapse process. In particular, we show that the model predicts a surface tension of $\\sim5 -15\\,$MeV$/{\\rm fm}^2$, rendering nucleation of quark matter possible during the early post-bounce stage of core collapse supernovae.  The LSMq, however, does not contain essential ingredients to describe nuclear matter, e.g. it does not reproduce features such as the saturation density and the binding energy. Therefore, the results obtained here should be considered with caution when applied to compact stars or the early universe. It is an effective theory for a first-order chiral phase transition in cold and dense strongly interacting matter, and allows for a clean calculation of the physical quantities that are relevant for homogeneous nucleation in the process of phase conversion. In the spirit of an effective model description, our results should be viewed as estimates that indicate that the surface tension is reasonably low and falls with baryon density, as one increases the supercompression. First-principle calculations in QCD in this domain are probably out of reach in the near future. Therefore, estimates within other effective models would be very welcome."
    },
    "1107/1107.1062_arXiv.txt": {
        "abstract": "We present the implementation of a halo based method for the reconstruction of the cosmic mass density field. The method employs the mass density distribution of dark matter haloes and its environments computed from cosmological N-body simulations and convolves it with a halo catalog to reconstruct the dark matter density field determined by the distribution of haloes. We applied the method to the group catalog of Yang \\etal (2007) built from the SDSS Data Release 4. As result we obtain reconstructions of the cosmic mass density field that are independent on any explicit assumption of bias. We describe in detail the implementation of the method, present a detailed characterization of the reconstructed density field (mean mass density distribution, correlation function and counts in cells) and the results of the classification of large scale environments (filaments, voids, peaks and sheets) in our reconstruction. Applications of the method include morphological studies of the galaxy population on large scales and the realization of constrained simulations. ",
        "introduction": "Galaxy surveys have become a useful mean to investigate the large scale structure of the Universe. Since late 60's with the Lick Galaxy Catalog (Shane \\& Wirtanen 1967), the large scale structure of the Universe has been studied through the distribution of galaxies in surveys like the CfA Redshift Survey (Huchra \\etal 1983), 2MASS Redshift Survey (Huchra 2000), the Las Campanas Redshift Survey (Shectman \\etal 1996), the 2dF survey (Colless \\etal 2001) and the Sloan Digital Sky Survey (York \\etal 2000). These surveys provide us with a wealth of data that have been used extensively in the study of the properties and evolution of galaxies as well as in studies of the properties of the mass distribution in the universe. Under the current paradigm of structure formation ($\\Lambda$CDM), galaxies are biased tracers of the mass distribution, which is thought to be dominated by dark matter. Nevertheless, observing the galaxy distribution is the only way we have to study the large scale cosmic mass distribution.  From the theoretical point of view, the cosmic mass distribution can be modeled as a smooth continuous function of the coordinates. From that assumption, important physical insight can be obtained about the process of structure formation and the dynamics of the Universe. However the observed galaxies represent discrete objects, therefore the discrete distribution of points has to be used to infer the smooth continuous mass density field behind the observed galaxy distribution. Such an investigation requires not only high amount of good quality data, but also appropriate procedures able to reliably reconstruct the smooth mass density field. These methods must be able to deal with the inherent difficulties associated with the observations like incompleteness, selection effects, geometrical constraints and redshift space distortions.  Fortunately, current galaxy redshift surveys provide the required amount and data quality to allow the reconstruction of the mass density field. Previously, many efforts have been made to reconstruct the density field using surveys like 1.2Jy IRAS redshift survey (Fisher \\etal 1995), 2dF Galaxy Redshift Survey (Erdo{\\u g}du \\etal 2004), 2MASS Redshift Survey (Erdo{\\u g}du \\etal 2006), the Luminous Red Galaxy Sample in the SDSS (Reid \\etal 2009), the 10k zCOSMOS survey (Kova{\\v c} \\etal 2010) and the SDSS galaxy redshift survey (Kitaura \\etal 2009, Jasche \\etal 2010). Almost of all of these works have based the reconstruction on the distribution of individual galaxies using the Wiener filter technique. Other methods have been proposed, like Delaunay tessellation (Schaap \\& van de Weygaert 2000), Monge-Amp{\\`e}re-Kantorovich method (Mohayaee \\etal 2006), Bayesian methods (Jasche \\etal 2010) and halo based methods (Wang \\etal 2009).  All of these methods use the galaxy distribution as tracers of the cosmic mass distribution, from them, they compute the cosmic mass distribution after smoothing the particle distribution on a grid and convolving the particle distribution with a given smoothing kernel. Since the reconstruction is based on the distribution of galaxies, and they may be biased tracers of the mass distribution, assumptions on the shape and functional dependence of the bias factor are needed (Erdo{\\u g}du \\etal 2006, Mo \\& White 1996).  Wang \\etal (2009) have proposed a halo based reconstruction method that uses dark matter haloes instead of galaxies as the point process behind the reconstruction procedure, while it simultaneously enables the inclusion of environments on the mass distribution around these haloes. Although they only tested their procedure against simulations, it seems to be quite promising, once one is willing to accept, first, the validity of the results of N-body simulations as true realizations of the large scale mass distribution, and second, assuming that from the observations one can build a reliable halo catalog to make the reconstruction. Currently both requirements seems to be fulfilled, given the good performance of the cosmological simulations in the concordance $\\Lambda$CDM model to reproduce the features of the cosmic large scale structure, and the current availability of data that makes possible to build catalogs of groups of galaxies in the local Universe (e.g. Tago \\etal 2008, Tago \\etal 2010, Yang \\etal 2007, Crook \\etal 2007, Berlind \\etal 2006). Particularly Yang \\etal (2007) provides a group catalog based on the identification of groups of galaxies that are located in the same dark matter halo. The realistic mass assignment in these groups provides a unique opportunity to refer to the dark matter haloes in the volume of the survey.  Reconstructions of the cosmic density field are necessary in the study of the large scale structure through galaxy surveys. As was already mentioned, they provide the continuous field that relates the observations with the theoretical model describing our understanding of the large scale structure. They are required to make calculations of the intrinsic properties of the density field, like power spectrum and mass variances (Jasche \\etal 2010, Tegmark \\etal 2004). Furthermore, they are used to extract information of the cosmological parameters and study the dynamics of the Universe (Percival \\etal 2010). Usually cosmic environments are defined as a function of the properties of the cosmic mass distribution (Forero-Romero \\etal 2009, Hahn \\etal 2007). Identification and classification of environments play a major role in the study of the physical mechanisms involved in the process of galaxy formation and evolution. High quality reconstructions of the local mass density field coupled with reconstructions of the velocity field can offer an interesting initial setup for the so called time machines that allow the realization of constrained cosmological simulations that can be used to study the time evolution of our local neighborhood (Kolatt \\etal 1996, Martinez-Vaquero \\etal 2007, 2009, Gottl\\\"{o}ber \\etal 2010, Lavaux \\etal 2010b, Nuza \\etal 2010).  In this work we address the problem of the reconstruction of the cosmic mass density field following the halo based technique presented in Wang \\etal (2009). We extend the method to observational data and apply it to the Fourth Data Release of the Sloan Digital Sky Survey (SDSS-DR4) using the group catalog built by Yang \\etal (2007) and the outputs of cosmological N-body simulations. We perform a series of tests to verify the quality and consistency of the reconstruction and make use of the results to classify the morphology of the large scale density field.  In this paper we first describe our methods, starting with the outline of the approach. Then we describe the construction of the halo catalog on which the reconstruction is based. Following, we show how to compute the typical mass distribution in and around dark matter haloes in simulations and to describe the procedures used to make the mass assignment in the reconstruction, and finally we put all steps together describing the procedure used to make the reconstruction. Then we show our results, discuss the properties of the reconstructed mass density field, compute different characteristics of the density field, and as an application, we perform the classification of the cosmic network from the reconstruction. ",
        "conclusions": "We have presented the implementation of a halo based reconstruction of the cosmic mass density field. The method combines the results of observations and simulations, convolving the mass density distribution in and around dark matter haloes computed from N-body simulations with the coordinates of haloes of a given mass identified in the group catalog of a galaxy survey. We have used the group catalog of Yang \\etal (2007) built from the SDSS-DR4 together with high resolution simulations of the formation of the structure in the standard spatially flat $\\Lambda$CDM model.  The reconstruction method produces a distribution of points sampling the cosmic mass density field that is determined by the mass and position of the dark matter haloes in the halo catalog, convolved with the estimated typical density distribution of mass in haloes and their environments. The final particle distribution resembles that obtained from cosmological simulations, and as the reconstruction method does not requires the definition of any scale length, it provides a high dynamical resolution. Tracer particles are more probably found in high density regions, however the adaptivity of the method allows to sample also the low density regions of the cosmic web.  Instead of imposing a smoothing length, the resolution and smoothness of the method is limited by two factors, the value of the mass threshold \\Mth~ and the mass of the sampling particles $m_{\\mathrm{rec}}$. The resolution of the reconstruction is set by the choice of the mass threshold \\Mth. Using lower values for the mass threshold produce a more detailed reconstruction of the density field at small scales. The selection of the mass threshold plays a role similar to the smoothing length in standard grid techniques of computing the density field, like cloud in cell (CIC), where structures below the threshold value (smaller than the smoothing length in the case of the CIC method) are erased and smoothed out in the density field. A key result is that, as was shown throughout the paper, the choice of the value of the mass threshold does not affect the global properties of the reconstructed density field.  The other parameter defining the the smoothness of the reconstruction is the mass of the particles used to sample the density field, $m_{\\mathrm{rec}}$. Low mass sampling particles produce a smooth and more well defined sampling of the density field, specifically in the case of haloes with masses close to the mass threshold \\Mth. In general $m_{\\mathrm{rec}}$ has to be lower than \\Mth, to allow the particle distribution to reliably reproduce the mass associated with haloes of masses close to the mass threshold.  It is worth to note that the reconstruction method is based on the concept of individual haloes, that is, main haloes in the language used in cosmological simulations. Due to observational constraints, this method would be better to be used in populations of haloes in a mass range above \\Mth$ > 10^{11}$\\hMsun, where is it is possible to reliably identify complete samples galaxies and groups of galaxies associated to haloes down to this mass value.  As it has been shown in section \\ref{sec:results}, the shape of the correlation function for our reconstructions using two different values of \\Mth~ converges to similar values at almost all scales. From this result and the fact that the method does not make any explicit assumption about biasing, the resulting reconstructions are a very suitable tool to address studies of galaxy properties and their relation with the large scale density field. The known effect of biasing, for haloes of different mass is introduced in a natural way in the reconstruction through the functions $\\eta(r,M)$.  A drawback of the method (and of any reconstruction method in general) is that it is strongly dependent on the properties of the catalog of haloes as well as on the details of the typical mass distribution in and around of haloes $\\eta(r,M)$ extracted from the simulations. The quality of the observations, and in particular, the properties of the halo catalog, have an influence on the quality of the reconstruction. It was shown that the mass assigned to haloes in the halo catalog gives values of the mean cosmic density that are off by a factor of around 2. Clearly better estimates of the masses of haloes may produce a better constraint. Despite the discrepancy, we see this result as motivating and as a good starting point in the improvement on the methods used to assign masses in the halo catalog like the one presented by Yang \\etal (2007). On the other hand, as we have shown, we can assume the masses in the halo catalog to be off by the same factor of $\\sim2$ as the density. This value is in complete agreement with the measured scatter for the mass assigned to haloes in the group catalog (Yang \\etal 2008, Wang \\etal 2008). We believe that the volume of the survey have no influence on this mismatch since it is totally independent on any assumption besides the cosmological parameters.  It is also true that assuming the typical density distribution of mass from simulations will make the method not appropriate to extract further cosmological information from the mass distribution in the halo catalog. Such an assumption is similar to assuming a power spectrum shape, as done in methods like the Wiener filter techniques. As it has been already mentioned, the power of the method relies on its physical simplicity, directly connected to the ideas of the halo model, and it is best suited to make studies of the properties of the mass distribution: large scale environments and environmental effects on galaxy formations, and problems where there is a pre-stablished cosmological model, like the realization of constrained simulations.  Besides the aforementioned sources of error coming from the observations, we have shown different potential sources of error in our reconstruction, and we have shown where the reconstruction is more strongly dominated by the data or by the priors of the method. Nevertheless is difficult to quantify the degree of error introduced in the reconstruction. We have estimated the amount of deviation between the expected and reconstructed mean mass density. This number may be taken partially as an estimation for the mean error in the reconstruction.  Analysis in the clustering pattern give also information about the degree of effects introduced by the method at different scales. But in general, as far as it has been shown in this and in previous implementations of the method, there is not yet a formal model to properly estimate the errors of the reconstructed density field in a given point, and will be part of further research to study in deep this issue.  To check for the consistency of our results we have made an exhaustive study of the properties of the reconstructed density field. We have shown that the reconstructed density field reproduces the mass correlation function. It clearly shows the two regions associated to the one halo and two halo terms. Comparing the correlation function for our reconstruction with the one for the halo catalog, we see that the amplitude of the correlation function is in agreement with the expected clustering at large scales, and by construction it reproduces the small scale clustering. More interesting is that it shows the bias-free nature of the method in the mass correlation function. We computed the properties of the mass distribution in spheres and verified the way the mass distribution follows a skewed, log-normal like distribution when computed at small scales. As expected (Coles \\& Jones 1991), this distribution broadens and shifts to $\\bar{\\delta}~\\sim 0$ for larger smoothing radii.  The classification of the cosmic web for the reconstructed density field also offered consistent results with the theoretical expectations. The volume filling fraction of the different characteristic environments (filaments, sheets, peaks and voids) in the reconstruction match the ones obtained from simulations. Only small differences within the expected cosmic variance are observed (Forero-Romero \\etal 2009).  Reconstructions of the density field as the one we have shown provide a way to map directly the cosmic mass distribution from the halo catalogs identified in the surveys without the necessity to assume any functional form for the bias factor. Provided that the data fulfills a minimum completeness requirements, coupling this method of reconstruction of the density field with a reconstruction of the velocity field will offer the opportunity of computing real space density fields. This will give us detailed information about the velocity field of the observed volume, information valuable to quantify and map the effects of the redshift space distortion, and usable to test the dynamics of the local neighborhood (Lavaux \\etal(2010a), Strauss \\& Willick 1995).  Furthermore, the results of the cosmic web classification obtained with the help of the reconstructed density field can be used to study the properties of the galaxies in the large scale environment, not only in their close environment (Tempel \\etal 2010, Weinmann \\etal 2009, 2008, Wang \\etal 2008, Park \\etal 2007).  The ability of the method to trace the mass distribution in an adaptive way, coupled with the ability to work in real space, makes the method a good choice to produce appropriate and high resolution density fields to be used as background for time machines used in the realization of constrained simulations structure formation (Kolatt \\etal 1996, Martinez-Vaquero \\etal 2009, Gottl\\\"{o}ber \\etal 2010, Lavaux \\etal 2010b, Nuza \\etal 2010)."
    },
    "1107/1107.5477_arXiv.txt": {
        "abstract": "We model the evolution of manganese relative to iron in the progenitor system of the globular cluster Omega Centauri by means of a self-consistent chemical evolution model. We use stellar yields that already reproduce the measurements of [Mn/Fe] versus [Fe/H] in Galactic field disc and halo stars, in Galactic bulge stars and in the Sagittarius dwarf spheroidal galaxy. We compare our model predictions to the Mn abundances measured in a sample of 10 red giant members and 6 subgiant members of $\\omega$\\,Cen. The low values of [Mn/Fe] observed in a few, metal-rich stars of the sample cannot be explained in the framework of our standard, homogeneous chemical evolution model. Introducing cooling flows that selectively bring to the cluster core only the ejecta from specific categories of stars does not help to heal the disagreement with the observations. The capture of field stars does not offer a viable explanation either. The observed spread in the data and the lowest [Mn/Fe] values could, in principle, be understood if the system experienced inhomogeneous chemical evolution. Such an eventuality is qualitatively discussed in the present paper. However, more measurements of Mn in $\\omega$\\,Cen stars are needed to settle the issue of Mn evolution in this cluster. ",
        "introduction": "\\label{sec:int}  The abundance ratios of chemical elements that are produced on different time scales by stars of different masses are essential probes of the history of chemical enrichment in different types of stellar populations. They tell which stars -- and in which proportions -- have contributed to the chemical evolution of a given system at any time. After Tinsley (1979), it has become customary to explain the trend of oxygen relative to iron as due to the different roles played by type Ia supernovae (SNeIa) and type II supernovae (SNeII) in the chemical enrichment of galaxies (see also Greggio \\& Renzini 1983; Matteucci \\& Greggio 1986). This interpretation is known as the `time-delay model'. Oxygen is produced mostly by core-collapse SNe on very short time scales, of the order of a few Myr to tenths of Myr. The bulk of iron, instead, comes later from SNIa explosions, on time scales ranging from 30~Myr to a Hubble time. Hence, the [O/Fe] ratio in a given system is high as long as SNeII dominate the interstellar medium (ISM) pollution, then decreases to solar and subsolar values when SNeIa become the major contributors to the production of Fe. The maximum SNIa rate depends on the adopted progenitor model, as well as on the assumed star formation history (Matteucci \\& Recchi 2001); thus, the measurements of [O/Fe] are a powerful diagnostics of the star formation history in galaxies.  Of particular interest are those elements whose yields depend on the metallicity of the parent stars; manganese (Mn) is one of them. McWilliam, Rich \\& Smecker-Hane (2003) confronted the [Mn/Fe] versus [Fe/H] relation in the Galactic bulge, in the solar neighbourhood and in the Sagittarius dwarf spheroidal galaxy (Sgr dSph). They suggested that the Mn yields from both SNeIa and SNeII are metallicity dependent. Their proposals were in agreement with extant nucleosynthesis calculations for massive stars (e.g. Arnett 1971; Woosley \\& Weaver 1995), but there was not modelling of explosive nucleosynthesis in SNeIa supporting their findings.  Shetrone et al. (2003) measured the Mn abundances of a dozen individual red giant stars in the Sculptor, Fornax, Carina and Leo~I dSphs. They also conclude that SNIa contributions to the synthesis of Mn must be metallicity dependent, with very little Mn produced until [Fe/H]~= $-$1.  Later on, Ohkubo et al. (2006) showed that, indeed, [Mn/Fe] -- and [Ni/Fe] -- in the ejecta of SNeIa depend on metallicity of SNIa progenitors. Furthermore, Cescutti et al. (2008) demonstrated, by means of self-consistent chemical evolution models, that the run of [Mn/Fe] with [Fe/H] in the three independent stellar systems -- the Galactic bulge, the solar neighbourhood and the Sgr dSph-- can be understood \\emph{only in terms of a metallicity-dependent yield of Mn} from SNeIa. The time-delay model alone is insufficient to explain the behaviour of [Mn/Fe] versus [Fe/H] in the three systems. Cescutti et al. (2008) propose that the yield of Mn in SNeIa increases with the metallicity of the progenitors, $Y_{\\rmn Mn} (Z) \\propto Z^{0.65}$. This is in agreement with the tight correlation between the Mn-to-Cr mass ratio in the ejecta of SNeIa and the metallicity of the progenitor found by Badenes, Bravo \\& Hughes (2008), $M_{\\rmn Mn}/M_{\\rmn Cr}$~= 5.3~$\\times$~$Z^{0.65}$.  More recently, Mn abundances have been studied for the first time on a significant metallicity range in the peculiar globular cluster $\\omega$\\,Centauri (Cunha et al. 2010; Pancino et al. 2011). In the metal-poor regime, Cunha et al. (2010) find that the LTE values of [Mn/Fe] in $\\omega$\\,Cen stars overlap those of their solar neighbourhood analogues. However, at variance with the solar neighbourhood trend, [Mn/Fe] declines in more metal-rich stars (see also Pancino et al. 2011). Non-LTE calculations confirm the conclusion of a well distinct pattern of [Mn/Fe] versus [Fe/H] in $\\omega$\\,Cen (Cunha et al. 2010). It is worth noticing that, in the metallicity range $-$2.7~$<$~[Fe/H] $< -$0.7, all other Galactic globulars have Mn abundances equivalent to those of halo field stars (Sobeck et al. 2006).  Although historically classified as a globular cluster, $\\omega$\\,Cen is possibly the naked nucleus of a small galactic satellite captured by the Milky Way many Gyr ago (e.g. Freeman 1993; Bekki \\& Freeman 2003). As such, it likely suffered a complex chemical enrichment history, marked by the occurrence of strong differential galactic winds (Romano et al. 2007, 2010a). The challenging bet is: can the low values of [Mn/Fe] observed in the more metal-rich stars of $\\omega$\\,Cen be explained in the context of its peculiar evolutive history?  In this paper, we contrast the evolution of Mn in the Milky Way with that in $\\omega$\\,Cen. The goal of this work is twofold. First, we want to add the analysis of another distinct stellar population to previous theoretical study of the evolution of Mn in different environments by Cescutti et al. (2008). Second, we want to get better insight into the mechanisms of formation and evolution of $\\omega$\\,Cen. The predictions of our models for the Milky Way and for $\\omega$\\,Cen are compared to both LTE and non-LTE Mn abundances in the two systems. The paper is organized as follows: in Sect.~2, we briefly review the relevant observational data. In Sect.~3, we describe the adopted chemical evolution models. In Sect.~4, we discuss the model results. In Sect.~5, we draw our conclusions. ",
        "conclusions": "In this paper we present the theoretical [Mn/Fe] versus [Fe/H] relationships predicted with our self-consistent chemical evolution model for $\\omega$\\,Cen, by using different prescriptions on Mn and Fe synthesis in stars. In particular, as for Mn production from SNeIa, we adopt either a metal-independent yield, by assuming Iwamoto et al.'s (1999) W7 model at all metallicities, or a metal-dependent one, by interpolating between models W70 and W7 of Iwamoto et al. (1999) or by adopting the empirical law of Cescutti et al. (2008). We also run models without Mn production from SNeIa. The theoretical relations are then compared to LTE and non-LTE data on Mn abundances in $\\omega$\\,Cen giants.  The adopted chemical evolution model reproduces all the major chemical properties of $\\omega$\\,Cen -- its metallicity distribution function, age-metallicity relation, average trends of several $\\alpha$-element-to-iron abundance ratios as functions of [Fe/H] (Romano et al. 2007). It also accounts for the presence of extreme He-rich stars (in the right percentage) and for the existence of a Na-O anticorrelation in the cluster (Romano et al. 2010a). However, the trend of decreasing [Mn/Fe] with increasing [Fe/H], displayed by both LTE and non-LTE data, is not reproduced by the model, which predicts instead a [Mn/Fe] ratio either increasing or constant in time, depending on the choice of stellar yields. The adoption of a metallicity-dependent, rather than metal-independent, Mn yield from SNeIa leads to an almost flat [Mn/Fe] versus [Fe/H] trend in $\\omega$\\,Cen only if the empirical law of Cescutti et al. (2008) is adopted. By interpolating between models W70 and W7 of Iwamoto et al. (1999), i.e. between models computed for metal-free and solar-metallicity SNeIa, respectively, we still produce a [Mn/Fe] ratio increasing with [Fe/H] in $\\omega$\\,Cen. On the basis of the discussion in Sect.~4 of the present paper and on previous results by Cescutti et al. (2008), we conclude that a flat trend has to be preferred over an increasing one.  In our chemical evolution model for $\\omega$\\,Cen, the chemical properties of the cluster are mainly driven by the action of strong differential galactic winds, which deeply affect the evolution of its dSph precursor (see Romano et al. 2007, and references therein, for details on the adopted scenario). Allowing for cooling flows or field star capture would not help to explain the presence of low-Mn stars in the cluster. Relaxing the hypothesis of homogeneous chemical evolution, i.e. allowing the stars to form from a mixture of ISM and individual SN ejecta, could eventually lead (for some stars) to the prediction of Mn abundances lower than predicted for the gas. However, this would hardly accommodate the Mn abundances of the two most metal-rich stars in Cunha et al.'s (2010) sample in a scenario of flat -- rather than decreasing -- Mn evolution with time.  We suggest that more measurements of Mn in $\\omega$\\,Cen stars at high metallicity are needed to finally set the issue of Mn evolution in $\\omega$\\,Cen. Interestingly, [Mn/Fe] in Sagittarius stays almost flat (Sbordone et al. 2007; see Fig.~\\ref{fig:all}) and Sagittarius is likely to be the local counterpart of the accretion episode which led to the formation of $\\omega$\\,Cen some 10 Gyr ago (Bellazzini et al. 2008; Georgiev et al. 2009; Carretta et al. 2010). Hence, it would be not surprising if future measurements of Mn in a much bigger sample of $\\omega$\\,Cen stars should reveal a constant run of Mn with [Fe/H], rather than the decreasing trend suggested by Cunha et al. (2010) on the basis of extant data. On the other hand, would future measurements confirm a fall of the [Mn/Fe] ratio towards higher metallicities, a revision of current scenarios of the formation of the cluster may be needed, since none of them is able to explain a decreasing trend of Mn in $\\omega$\\,Cen. We have shown that, if Mn production from SNeIa is totally suppressed in $\\omega$\\,Cen, a [Mn/Fe] ratio decreasing with time (metallicity) can be found in the cluster. While it is pointless to speculate on this theoretical result until more Mn data for $\\omega$\\,Cen stars become available, we note that Jonhson \\& Pilachowski (2010) observe consistently elevated [$\\alpha$/Fe] ratios for nearly all stars in the cluster (their sample totals 855 giants) and interpret this as evidence against a significant contribution to $\\omega$\\,Cen's chemical enrichment from SNeIa."
    },
    "1107/1107.5688_arXiv.txt": {
        "abstract": "{} { We carried out high-resolution spectroscopy and {\\it BV(I)}$_C$ photometric monitoring of the two fastest late-type rotators in the nearby $\\beta$ Pictoris moving group, HD~199143 (F7V) and CD$-$64$^{\\circ}$1208 (K7V). The motivation for this work is to investigate the rotation periods and photospheric spot patterns of these very young stars, with a longer term view to probing the evolution of rotation and magnetic activity during the early phases of main-sequence evolution. We also aim to derive information on key physical parameters, such as rotational velocity and rotation period.} { We applied maximum entropy (ME) and Tikhonov regularizing (TR) criteria to derive the surface spot map distributions of the optical modulation observed in HD~199143 (F7~V) and CD$-$64$^{\\circ}$1208 (K7~V). We also used cross-correlation techniques to determine stellar parameters such as radial velocities and rotational velocities. Lomb-Scargle periodograms were used to obtain the rotational periods from differential magnitude time series.} { We find periods and inclinations of 0.356 days and 21.5~deg for HD~199143, and 0.355 days and 50.1~deg for CD$-$64$^{\\circ}$1208. The spot maps of HD~199143 obtained from the ME and TR methods are very similar, although the latter gives a smoother distribution of the filling factor.  Maps obtained at two different epochs three weeks apart show a remarkable increase in spot coverage amounting to $\\sim 7$\\% of the surface of the photosphere over a time period of only $\\sim 20$ days. The spot maps of CD$-$64$^{\\circ}$1208 from the two methods show good longitudinal agreement, whereas the latitude range of the spots is extended to cover the whole visible hemisphere in the TR map. The distributions obtained from the first light curve of HD~199143 show the presence of an extended and asymmetric active longitude with the maximum filling factor at longitude $\\sim 325^{\\circ}$. A secondary active longitude is present at $\\sim 100^{\\circ}$. The spotted area distributions on CD$-$64$^{\\circ}$1208 show two active longitudes separated by about $180^{\\circ}$, which is not unusual on such very active stars.  } {} ",
        "introduction": "Moving groups (MGs) are kinematically coherent groups of stars that probably share a common origin---the evaporation of an open cluster, the remnants of a star formation region, or a juxtaposition of several small star formation bursts at different epochs in adjacent cells of the velocity field. Recent years have seen the discovery of several MGs within 100\\,pc of the Sun (e.g.; TW Hydrae, Horologium-Tucana and $\\beta$ Pictoris). Stars in these groups range between millions and tens of million years in age, and their study should tell us about the formation and evolution of sparse stellar associations. The stars themselves provide a rare glimpse of the very earliest main-sequence evolution and an opportunity to study magnetic activity at its extremes. Young late-type fast rotating members of MGs offer a unique laboratory for studying stellar dynamos, magnetic structures and coronal heating.  The $\\beta$ Pictoris moving group (BPMG) is one of the youngest \\citep[$\\sim$12\\,Myr,][]{kai04} and the closest MGs to the Sun \\citep[$\\sim$36\\,pc,][]{zuc01}. \\cite{zuc01} and \\cite{son03} report a total of 35 known members of the BPMG, each with one or more characteristics indicative of extreme youth. \\cite{kai04} found a high lithium abundance in one of the BPMG members.  The best-known members of the BPMG are early-type stars, whereas few studies have concentrated on the late-type stars representing three fourths of the known BPMG membership. The rotational properties of the late-type members have been recently investigated by \\cite{messina10} as part of the RACE-OC project.  We have carried out high resolution spectroscopy and {\\it BV(I)$_C$} photometric monitoring of a selected sample of young late-type stars (spectral type later than F2). The sample chosen for this work has been selected from previously established members of the BPMG, based on photometric and kinematic properties. The motivation for this work is to investigate the rotation periods and photospheric spot patterns of these very young stars, with a longer-term view to probing the evolution of rotation and magnetic activity during the early phases of main-sequence evolution.  Here, we present results for two members of our sample, HD~199143 and CD$-$64$^{\\circ}$1208. These are the fastest rotators among the known late-type stars in the BPMG. HD~199143 has a $v\\sin i$ of 128 km\\,$s^{-1}$ \\citep{tor06} which is the largest of any known single solar-like star within 50~pc. \\cite{anc00} proposed that HD 199143 has been spun up by accretion of material from a close T Tauri-like companion responsible for the emission lines, the ultraviolet variability and the excess infrared emission. In the case of CD$-$64$^{\\circ}$1208, $v \\sin i > 100 $~km~s$^{-1}$ \\citep{zuc01} and its unusually strong Li for a 10 Myr old star has been postulated as possibly being due to its very fast rotation \\citep{sod93}. Both stars represent a special opportunity to study the stellar atmospheric activity in young solar-like ultra-fast rotators.  { Further information on the targets and observations are given in Sect.~\\ref{sec:targets} and Sect.~\\ref{sec:obs} respectively. Stellar parameters, time series analysis and phasing of the light curves are described in Sect.~\\ref{sec:parameters}, Sect.~\\ref{sec:timeseries} and Sect.~\\ref{sec:phasinglightcurves} respectively. The spot map distributions are shown in Sect.4~ref{sec:spotmap}, while in Sect.~\\ref{sec:results} we present the results. The conclusions are given in Sect.~\\ref{sec:conclusions}.} ",
        "conclusions": "\\label{sec:conclusions}  We have analysed photometric and high resolution spectroscopic observations, made during the period 2004--2005, to investigate the surface activity of the two fastest rotating late-type stars known in the $\\beta$ Pictoris young moving group.  In addition to surface spot map distributions, these observations have yielded information on the key physical parameters of rotational velocity, rotation period, inclinations angle and radial velocity.  \\begin{enumerate}  \\item Radial velocity standard stars have been used to calculate accurate heliocentric radial velocities for each observation by using the cross-correlation technique. The mean velocities for HD~199143 and CD$-$64$^{\\circ}$1208 are 2.2 $\\pm$ 2.9 km~s$^{\\rm -1}$ and 2.5 $\\pm$ 2.2 km~$s^{\\rm -1}$, respectively. These values were used to calculate the Galactic space-velocity components ($U$,$V$,$W$) and confirmed their membership of the BPMG.  \\item From cross-correlation analysis of our spectra, projected rotational velocities, $v \\sin i$, of 115.5$\\pm$7.5~km~s$^{-1}$ and 121.3$\\pm$15.3~km~s$^{-1}$ have been measured for HD~199143 and CD$-$64$^{\\circ}$1208, respectively.  \\item Rotation periods of 0.356$\\pm$0.004 days for HD~199143 and 0.355$\\pm$0.040 days for CD$-$64$^{\\circ}$1208 have been derived from the time series of V-band differential magnitudes.  \\item The inclination, $i$, of the stellar rotation axis with respect to the line of sight has been derived from the observed $v \\sin i$, the estimated stellar equatorial radius $a$, and the rotation period $P_{\\rm rot}$. We find inclination angles of $21\\fdg 5$ and $50^{\\circ}$ with an uncertainty of $10^{\\circ}-15^{\\circ}$ for HD~199143 and CD$-$64$^{\\circ}$1208, respectively.  \\item The spot maps of HD~199143 obtained using the TR and ME criteria are remarkably similar; the TR criterion gives a smoother distribution of the filling factor, as expected. The latitude of the spots is confined below $\\sim 60^{\\circ}$ in the first map and below $\\sim 30^{\\circ}$ in the second map. Light curve amplitude variations over the different epochs of observation require an increase of $\\sim 7$\\% in the spot filling factor, unevenly distributed in longitude, on a time scale of only $\\sim 20$ days.  \\item The spot maps of CD$-$64$^{\\circ}$1208 show good longitudinal agreement whereas the latitude range of the spots is extended to cover the whole visible hemisphere in the TR map.  \\item The distributions obtained from the first light curve of HD~199143 show the presence of an extended and asymmetric active longitude with the maximum filling factor at longitude $\\sim 325^{\\circ}$. A secondary active longitude is present at $\\sim 100^{\\circ}$.  It appears to have grown remarkably in the second light curve maps, while the primary active longitude has migrated slightly towards decreasing longitudes. The spotted area distributions on CD$-$64$^{\\circ}$1208 show two active longitudes separated by about $180^{\\circ}$ which is not unusual on such very active stars. \\end{enumerate}"
    },
    "1107/1107.0314_arXiv.txt": {
        "abstract": "{This study is part of a large project to study the physics of accretion and molecular outflows towards a selected sample of high-mass star-forming regions that show evidence of infall and rotation from previous studies.} {We wish to make a thorough study at high-angular resolution of the structure and kinematics of the HMCs and corresponding molecular outflows in the high-mass star-forming region G24.78+0.08.} {We carried out SMA and IRAM PdBI observations at 1.3 and 1.4~mm, respectively, of dust and of typical high-density and molecular outflow tracers with resolutions of $<1\"$. Complementary IRAM 30-m $^{12}$CO and $^{13}$CO observations were carried out to recover the short spacing information of the molecular outflows.} {The millimeter continuum emission towards cores G24 A1 and A2 has been resolved into 3 and 2 cores, respectively, and  named A1, A1b, A1c, A2, and A2b. All these cores are aligned in a southeast-northwest direction coincident with that of the molecular outflows detected in the region, which suggests a preferential direction for star formation in this region. The masses of the cores range from 7 to 22~$M_\\odot$, and the rotational temperatures from 128 to 180~K. The high-density tracers have revealed the existence of 2 velocity components towards A1, one of them peaks  close to the position of the millimeter continuum peak and of the \\HC\\ region, and is associated  with the velocity gradient seen in \\MCN\\ towards this core, while the other one peaks southwest of core A1 and is not associated with any millimeter continuum emission peak. The position-velocity plots along outflow A and the $^{13}$CO~(\\jdu) averaged blueshifted and redshifted emission indicate that this outflow is driven by core A2. Core A1 apparently is not driving any outflow. The knotty appearance of the highly collimated outflow C and the $^{12}$CO position-velocity plot suggest an episodic outflow, where the knots are made of swept-up ambient gas.} {} ",
        "introduction": "Molecular outflows, infall and rotation are ubiquitous phenomena during the earliest stages of the formation of stars of all masses and luminosities. During the last few years, much effort has been devoted to studying and describing the physical properties of embedded massive protostars and their molecular outflows (e.g., Shepherd \\& Churchwell 1996a,b; Zhang et al.~\\cite{zhang01}; Beuther et al.~\\cite{beuther02}; L\\'opez-Sepulcre~\\cite{lopez09}). The presence of rotating structures around massive young stars has also been asserted thanks to the detection of velocity gradients perpendicular to the axis of the molecular outflows in the cores (see Cesaroni~\\cite{cesa07} and reference therein). These studies have revealed two different kinds of structures surrounding B- and O-type stars. On the one hand, B-type stars are surrounded by centrifugally supported disks, as seen in the prototypical B-type protostar IRAS~20126+4104 (Cesaroni et al.~\\cite{cesa05}; Keto \\& Zhang~\\cite{keto10}). On the other hand, despite the fact that theory predicts the existence of both  large rotating structures (e.g.\\ Peters et al.~\\cite{peters10}) and accretion disks (e.g.\\ Krumholz et al.~\\cite{krumholz09}) around massive O-type stars, such disks have been elusive to date. One finds, instead, huge ($\\sim$0.1 pc), massive (a few 100~$M_\\odot$) rotating toroids, which are non Keplerian and very likely non-equilibrium structures (Beltr\\'an et al.~\\cite{beltran11}). Finally,  infall onto young O-type stars has been measured through high-angular resolution observations of lines showing redshifted self-absorption or absorption against a background bright continuum source (Keto et al.~\\cite{keto87}; Zhang et al.~\\cite{zhang98}; Beltr\\'an et al.~\\cite{beltran06}), or by measuring the velocity field of the ionized gas surrounding a newly formed massive star (see Keto~\\cite{keto02}).  \\begin{table*} \\caption[] {Parameters of the maps} \\label{par_obs} \\begin{tabular}{lcccccc} \\hline &&&\\multicolumn{2}{c}{Synthesized beam$^{\\rm a}$} \\\\ \\cline{4-5} &&\\multicolumn{1}{c}{Frequency} & \\multicolumn{1}{c}{$HPBW$} & \\multicolumn{1}{c}{P.A.}& \\multicolumn{1}{c}{Spectral resolution}& \\multicolumn{1}{c}{rms noise$^{\\rm b}$} \\\\ \\multicolumn{1}{c}{Observation} & \\multicolumn{1}{c}{Telescope$^{\\rm a}$}& \\multicolumn{1}{c}{(GHz)} & \\multicolumn{1}{c}{(arcsec)} & \\multicolumn{1}{c}{(deg)}  & \\multicolumn{1}{c}{(\\kms)} & \\multicolumn{1}{c}{(mJy~beam$^{-1}$)} \\\\ \\hline continuum                                &SMA &225.31 &$0.95\\times0.75^{\\rm c}$  &61$^{\\rm c}$ &-- &3.0$^{\\rm c}$  \\\\ continuum                                &PdBI &214.30 &$0.89\\times0.29$ &10 &-- &3.2 \\\\ $^{13}$CH$_3$CN (12$_K$--11$_K$)         &PdBI &214.374$^{\\rm d}$ &$0.89\\times0.29^{\\rm e}$&10 &0.25 &15\\\\ C$^{18}$O (\\jdu)                         &SMA &219.560 &$0.83\\times0.64^{\\rm e}$ &55  &0.6 &35 \\\\ HNCO (10$_{0,10}$--9$_{0,9}$)            &SMA &219.798 &$0.72\\times0.51$ &46 &0.6 &45 \\\\ H$_2$$^{13}$CO (3$_{1,2}$--2$_{1,1}$)    &SMA &219.909 &$0.72\\times0.51$ &46 &0.6 &45 \\\\ SO (6$_5$--5$_4$)                        &SMA &219.949 &$0.83\\times0.64^{\\rm e}$ &55 &0.6 &40 \\\\ CH$_3$OH (8$_{0,8}$--7$_{1,6}$)          &SMA &220.079 &$0.66\\times0.48$ &48 &0.6 &45 \\\\ HCOOCH$_3$ (17$_{2,15}$--16$_{4,2}$)     &SMA &220.167 &$0.71\\times0.50$ &45 &0.6 &45 \\\\ CH$_2$CO (1$_{1,11}$--10$_{1,10}$)       &SMA &220.178 &$0.71\\times0.50$ &45 &0.6 &45 \\\\ HCOOCH$_3$ (17$_{4,13}$--16$_{4,12}$)    &SMA &220.190 &$0.71\\times0.50$ &45 &0.6 &45 \\\\ $^{13}$CO (\\jdu)                         &SMA &220.399 &$0.83\\times0.64^{\\rm e}$ &55 &1.0 &35 \\\\ $^{13}$CO (\\jdu)                         &SMA+IRAM~30-m &220.399 &$3.57\\times2.02$ &71 &1.0 &116 \\\\ \\MCNII\\ (12$_K$--11$_K$)                 &SMA &220.638$^{\\rm d}$ &$0.67\\times0.47$ &47 &0.6 &45\\\\ C$_2$H$_5$CN (25$_{2,24}$--24$_{2,23}$)  &SMA &220.661 &$0.67\\times0.48$ &47 &0.6 &45 \\\\ \\MCN\\ (12$_K$--11$_K$)                   &SMA &220.747$^{\\rm d}$ &$0.71\\times0.50$ &46 &0.6 &40\\\\ \\MCN\\ $v_8=1$ ( $K,l$=6,$+1$)            &SMA &221.312  &$0.66\\times0.47$ &46 &0.6 &40 \\\\ \\MCN\\ $v_8=1$ ($K,l$=3,$-1$)             &SMA &221.338 &$0.66\\times0.47$ &46 &0.6 &40 \\\\ CH$_3$OH (8$_{1,8}$--7$_{0,7}$)E         &SMA &229.759 &$0.65\\times0.47$ &48 &0.6 &50 \\\\ CH$_3$OH (19$_{5,14}$--20$_{4,17}$)A$-$  &SMA &229.939 &$0.65\\times0.47$ &47 &0.6 &45 \\\\ CH$_3$OH (3$_{2,2}$--4$_{1,4}$)E         &SMA &230.027 &$0.65\\times0.47$ &48 &0.6 &45 \\\\ $^{12}$CO (\\jdu)                         &SMA &230.538 &$0.86\\times0.68^{\\rm e}$ &54 &1.5 &40 \\\\ $^{12}$CO (\\jdu)                         &SMA+IRAM~30-m &230.538 &$3.37\\times1.96$ &70 &1.0 &116 \\\\ OCS (19--18)                             &SMA &231.061 &$0.69\\times0.50$ &46 &0.6 &50 \\\\ $^{13}$CS (5--4)                         &SMA &231.221 &$0.70\\times0.49$ &46 &0.6 &50 \\\\ \\hline \\end{tabular}  $^a$ The synthesized CLEANed beams for maps with the ROBUST parameter of Briggs~(\\cite{briggs95}) set equal to 0.\\\\ $^b$ For the molecular line observations the 1$\\sigma$ noise is per channel.\\\\ $^c$ The synthesized beam for the SMA VEX configuration map is $0\\farcs55\\times0\\farcs40$ at P.A.=52$\\degr$, and the rms noise is 3.2~mJy~beam$^{-1}$. \\\\ $^d$ The frequency is that of the $K$=0 component. \\\\ $^e$ The synthesized CLEANed beams for maps using natural weighting. \\end{table*}     G24.78+0.08 (hereafter G24) is one of the most studied massive star-forming regions. The simultaneous presence of the three elements fundamental for star formation (outflow, infall, and rotation) has been reported in one of the massive objects embedded in this clump. This region, which is located at a distance of 7.7~kpc, has been studied in great detail by our group (Codella et al.~\\cite{codella97}; Furuya et al.~\\cite{furuya02}; Cesaroni et al.~\\cite{cesa03}; Beltr\\'an et al.~\\cite{beltran04}, Beltr\\'an et al.~\\cite{beltran05}, Beltr\\'an et al.~\\cite{beltran06}, Beltr\\'an et al.~\\cite{beltran07}; Moscadelli et al.~\\cite{mosca07}; Vig et al.~\\cite{vig08}). G24 contains a cluster of Young Stellar Objects (YSOs), three of which have rotating toroids associated with them (Beltr\\'an et al.~\\cite{beltran04}). These rotating toroids seem to be associated with two compact molecular outflows, first mapped in CO by  Furuya et al.~(\\cite{furuya02}). In one of these rotating cores, G24~A1, the presence of an early-type star is witnessed by an embedded hypercompact (HC) \\HII region, first detected at 1.3~cm by Codella et al.~(\\cite{codella97}) and later studied at very high-angular resolution at 1.3 and 0.7~cm by Beltr\\'an et al.~(\\cite{beltran07}). The properties of the ionized gas have also been studied through continuum emission and recombination lines by Galv\\'an-Madrid et al.~(\\cite{galvan08}) and Longmore et al.~(\\cite{longmore09}). This \\HC\\ region, which has a diameter of $\\la$$0\\farcs2$ ($\\la$1,500~AU), is located at the geometrical center of the toroid. The free-free continuum spectrum resembles that of a classical (Str\\\"omgren) \\HII region around a zero-age main sequence star of spectral type O9.5, corresponding to a mass of $\\sim 20~M_\\odot$. Beltr\\'an et al.~(\\cite{beltran06}) mapped the core G24~A1 in NH$_3$ and detected redshifted absorption, indicating that the toroid is undergoing infall towards the \\HC\\ region.    We have carried out new 1.3~mm interferometric observations of this region at sub-arcsecond resolution with the Submillimeter Array (SMA)\\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.} (Ho et al.~\\cite{ho04}) of several outflow and high-density tracers. The goals of this study were to better characterize the structure and kinematics of the Hot Molecular Cores (HMCs) and molecular outflows in this region, and in particular, to identify the powering source of the outflow observed towards G24~A1. In order to be sensitive to extended structures filtered out by the interferometer, we have also mapped the region over $2'$ with the IRAM 30-m telescope\\footnote{IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).}. In addition, we used the enhanced capabilities of the IRAM Plateau de Bure (PdBI) to observe this region in \\MCNI~(\\jdo) at very high-angular resolution to investigate whether the massive toroids were hiding true (Keplerian) circumstellar disks in their interiors. Unfortunately, the angular resolution of the observations has not been enough to detect circumstellar disks embedded in the massive toroids, but it has allowed us to study the velocity structure of the cores with unprecedented resolution. The observational details are given in Sect.~2, while the results are illustrated in Sect.~3 and analyzed and discussed in Sect.~4. Finally, the conclusions are drawn in Sect.~5.    \\begin{figure*} \\centerline{\\includegraphics[angle=-90,width=16cm]{cont_iram_sma.eps}} \\caption{({\\it Left panel}) SMA map of the 1.3~mm continuum emission obtained combining both compact and very extended configurations.  Contour levels are 3, 6, 9, 12, 15, 35, and 60 times  $\\sigma$, where 1$\\sigma$ is 3\\mjy.  The crosses mark the positions of the sources identified by Furuya et al.~(\\cite{furuya02}) and Beltr\\'an et al.~(\\cite{beltran04}). The  triangle marks the position of the \\HC\\ region in G24 A1 (Beltr\\'an et al.~\\cite{beltran07}). The SMA synthesized beam is shown in the lower left corner. ({\\it Right panel}) Very extended configuration only SMA map at 1.3~mm ({\\it contours}) overlaid on the PdBI map at 1.4~mm ({\\it grayscale}) towards G24 A1, A2 and D. Contour and grayscale levels are 3, 6, 9, 12, 15, and 20 times $\\sigma$, where 1$\\sigma$ is 3.2\\mjy. The SMA synthesized beam is shown in the lower left corner and the PdBI one in the lower right corner. } \\label{g_cont} \\end{figure*} ",
        "conclusions": "We analyzed millimeter data obtained with the SMA and the IRAM PdBI interferometers, and the IRAM 30-m telescope of the dust and gas emission towards the cores in the high-mass star-forming region G24.78+0.08. The synthesized beam of $<1''$ of the observations has allowed us to study with unprecedented high-angular resolution the structure and the velocity field of the cores, and the molecular outflows powered by the YSOs embedded in the HMCs.  The millimeter continuum emission towards the cores A1 and A2 has been resolved into additional cores. Core A1 has been resolved into 3 cores named A1, A1b, and A1c, whereas core A2 has been resolved into 2 named A2 and A2b. Interestingly, the cores are aligned in a southeast-northwest direction coincident with that of the molecular outflows detected in the region. This suggests a preferential direction for star formation in this region that could possibly depend on the direction of the magnetic field. It could also indicate star formation in filaments or sheets. The core associated with the \\HC\\ region is A1. The masses of the cores for which it has been possible to estimate the rotational temperature range from 7 to 22~$M_\\odot$. The rotational temperatures range from 128 to 180~K.  The velocity gradients towards cores A1 and A2, first detected by Beltr\\'an et al.~(\\cite{beltran04}) in \\MCN~(\\jdo), have been confirmed by the new SMA observations. Unfortunately, the sensitivity of the observations have not allowed us to  better constrain the direction of the velocity gradient detected in CS~(3--2) by Beltr\\'an et al.~(\\cite{beltran04}).  The \\MCNII, \\MCNI\\ and C$_2$H$_5$CN observations have revealed the existence of 2 velocity components towards A1. The component at $\\sim$111~\\kms\\ is associated with the velocity gradient seen in \\MCN\\ and peaks close to the position of the millimeter continuum peak and of the \\HC\\ region. The  component at $\\sim$106~\\kms\\ peaks towards the SW of core A1 and does not seem to be associated with a millimeter continuum emission peak.  The two molecular outflows in the region have been mapped in different tracers. The position-velocity plots along outflow A and the $^{13}$CO~(\\jdu) averaged blueshifted and redshifted emission indicate that this outflow is driven by core A2. Regarding core A1, the position-velocity plots seem to indicate that this core is not powering any outflow. However, due to the fact that the outflow structure is very complex towards cores A1 and A2, we cannot totally discard the possibility that core A1 could be powering an additional outflow in the region. The outflow C is highly collimated and clumpy. The clumpiness and the $^{12}$CO~(\\jdu) position-velocity plot suggest that this outflow could be episodic, and that the outflow clumps are made of swept-up ambient gas."
    },
    "1107/1107.4317_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\noindent Electroweak baryogenesis (EWBG) ~\\cite{baryogenesis,reviews} is a very elegant mechanism for generating the baryon asymmetry of the Universe (BAU). It relies on physics at the weak scale and can therefore be tested at present accelerator energies, in particular at the LHC. During the electroweak phase transition (EWPT) bubbles of the broken phase are nucleated and expand. Particles in the plasma are reflected off the bubble walls where CP is violated and CP-violating currents may be generated. If the currents efficiently diffuse into the unbroken phase they may be converted into a baryon asymmetry by the action of the baryon number violating sphaleron processes~\\cite{sphalerons}. The baryon asymmetry then flows into the interior of the bubble where it is preserved provided that the sphaleron interactions are sufficiently switched off in the broken phase which defines a strong enough first order phase transition.   In this work we point out an effect that seems to have passed unnoticed so far and that may require a stronger first order phase transition, for EWBG purposes. During the first order phase transition magnetic fields are unavoidably generated~\\cite{ereview}.  Bubble collisions generate a level of turbulence and hence vorticity in the fluid. The turbulent conducting fluid develops magnetic turbulence resulting in magnetic fields on all scale sizes. The turbulence in the fluid amplifies whatever seed fields are present to finite-amplitude large-scale size magnetic fields \\cite{Stevens}.  The relevant time scale for the amplification of fields on length scale $\\ell$ is of order $(\\ell/R_{\\rm b})t_{\\rm pt}$, where $R_{\\rm b}$ is the radius of the bubble moving with velocity $v_{\\rm w}$ and $t_{\\rm pt}\\sim R_{\\rm b}/v_{\\rm w}$ is the duration of the phase transition. If the field growth is exponential fields on scales $\\ell\\lesssim R_{\\rm b}$ can be amplified by many $e$-folds~\\cite{baym,sigl}. When the magnetic turbulence becomes fully developed the kinetic energy of the turbulent flow is equipartitioned with that of the magnetic field energy implying that the magnetic fields $B(R_{\\rm b})$ on the size of the bubble radius is $B^2/2={\\cal O}(v_{\\rm f}^2)\\rho_\\gamma$, where $\\rho_\\gamma\\simeq (\\pi^2 g_*/30) T^4$ is the energy density of the electroweak plasma carrying $g_*$ relativistic degrees of freedom, $T$ is the temperature and $v_{\\rm f}$ is the fluid velocity~\\cite{baym,sigl}. This means that a magnetic field of size \\begin{equation} \\label{a} B\\sim 0.4\\left(\\frac{v_{\\rm f}}{0.05}\\right)\\,  T^2\\, , \\end{equation} can be generated via turbulence.  One possible mechanism for generation of magnetic seed fields is by the dipole electromagnetic charge layer that develops on the surfaces of the bubbles as a consequence of baryon asymmetry~\\cite{baym,sigl}. The rotation of the dipole charge layer thus sets up a current in the fluid.  The magnetic field may be generated from these currents in the bubble walls and then be amplified to the equipartition value (\\ref{a}) by exchange of energy with the turbulent fluid. Higgs phase gradients can also act as a source for gauge fields at bubble collisions~\\cite{kibble,ae}.  Of course magnetic fields from the electroweak transition can survive only on scales on which magnetic diffusion has not had time to wash out the field correlations. This means that magnetic fields on the bubble radius scale $R_{\\rm b}$ dies off on a time scale $\\sim \\sigma R_{\\rm b}^2$, where $\\sigma\\sim 10\\, T$~\\cite{con} is the conductivity of the plasma. Being bubble sizes ${\\cal O}(10^{-2}-10^{-3})$ of the Hubble radius at the electroweak phase transition magnetic fields set up at the electroweak phase transition die away on time scales much larger than the Hubble time at the phase transition.  Lacking a more precise calculation about the exact magnitude of the generated magnetic field at the electroweak phase transition we will parametrize it through the dimensionless parameter $b$ as \\begin{equation} B=b\\,T^2\\, , \\label{b} \\end{equation} while from what we have just discussed above values of $b\\lesssim 0.4$ seem quite plausible. Such values are  comfortably smaller than those deduced by imposing that the energy density stored in the magnetic field at the characteristic scale of production (in our scale the bubble radius) does not appreciably alter the  dynamics of primordial nucleosynthesis and the structure of the CMB anisotropies \\cite{caprini}.   Now the key point is that sphaleron configurations do possess a magnetic dipole moment~\\cite{Manton}.  In the background of a magnetic field in the bubble of the broken phase the coupling with the dipole moment lowers the height of the sphaleron barrier so that thermal fluctuations are more effective in producing topological transitions~\\cite{CGPR}.  Therefore in order to preserve the baryon asymmetry within the bubble where a magnetic field is produced by the phase transition it is then necessary to require a stronger first order phase transition than in the case in which one neglects the presence of the magnetic field.  The previous comments do apply to any first order phase transition generating EWBG. However it has been shown that EWBG cannot be realized within the Standard Model (SM) framework~\\cite{AndH, improvement, twoloop, nonpert,CPSM} and it is neither feasible in its Minimal Supersymmetric extension (MSSM) for arbitrary values of its parameters~\\cite{early, mariano1, mariano2}. A particular region in the space of supersymmetric mass parameters was found in the MSSM, dubbed under the name of light stop scenario (LSS)~\\cite{CQW, Delepine, CK, FL, JoseR, JRB, Carena:1997gx, Carena:1997ki, CJK, Iiro2, Toni2, Toni3, Worah, Schmidt, Cline:2000kb, Carena:2000id, Konstandin:2005cd, Cirigliano:2006dg, qn, chung1, chung2, chung3, window}, where EWBG had the potential of being successful. In particular the condition that sphaleron interactions are inhibited in the broken phase required a sufficiently strong first order phase transition imposing absolute upper bounds on the lightest CP-even Higgs and right-handed stop masses, $m_H\\lesssim 127$ GeV and $m_{{\\widetilde t}_R}\\lesssim 120$ GeV~\\cite{window}. This is the so-called MSSM baryogenesis window.  In this paper we will consider the EWPT in the MSSM and re-analyze the EWBG constraints taking into account the magnetic field produced by the phase transition bubbles. It is then clear that the presence of a magnetic field inside the bubble combined with a non-vanishing magnetic dipole moment of the sphaleron will lower the upper bound on the lightest CP-even Higgs mass and might close the present EWBG window. However as we have previously stated in the absence of a precise enough calculation we will parametrize the magnetic field by the dimensionless parameter in Eq.~(\\ref{b}) so that the results in this paper could be interpreted as upper bounds on the magnitude of the parameter $b$. In this way the MSSM EWBG scenario can only be disproved by future experimental results on the Higgs and right-handed stop masses and/or a more precise theoretical calculation of the produced magnetic field.  The paper is organized as follows. In Section~\\ref{sec:EWBG} we provide a short summary of the LSS, while in Section~\\ref{sec:sphaleron} we briefly discuss the energy of the sphaleron in a magnetic field, deferring the technical details to Appendices~\\ref{appendice1} and \\ref{appendice2}.  Section~\\ref{sec:results} contains our numerical results and  Section~\\ref{sec:conclusions} the conclusions. ",
        "conclusions": "\\label{sec:conclusions}   In this paper we have pointed out that the sphaleron magnetic dipole moment couples to the magnetic field generated during a first order EWPT. This has the effect of lowering the sphaleron energy in the broken phase and consequently the baryon asymmetry is washed out more easily.  We have not attempted to compute precisely the magnetic field generated at the phase transition,  but rather we have considered it as a free parameter within a plausible range.  In particular we have focused on the MSSM in the most favourable situation for EWBG, the light stop scenario, and explored the dependence of the sphaleron energy on the parameters of the model and on the magnetic field.  We have shown that it is possible to have the baryon asymmetry preserved even in presence of a magnetic field by lowering the Higgs mass and/or the stop mass, or increasing the soft scalar mass $\\widetilde m$. However even for moderate values of the magnetic field the required Higgs mass can fall below the present experimental bound and the window for EWBG in the MSSM gets closed.  The main conclusions of this paper are twofold.  On the one hand our calculation cries out  for an accurate determination of $B$ and its evolution during the phase transition in order to confirm if the magnetic field produced by the electroweak phase transition endangers the EWBG in the LSS, as our  analysis seems to indicate. On the other hand, the magnetic field produced by phase transition bubbles  threatens any model where the BAU is generated by a first order phase transition. In principle, any model of EWBG where the phase transition (evaluated at zero magnetic field) is never extremely strong should be jeopardized by the presence of a magnetic field."
    },
    "1107/1107.4856_arXiv.txt": {
        "abstract": "{The \\astrobj{Ursa Major group} is a nearby stellar supercluster which, while not gravitationally bound, is defined by co-moving members. \\astrobj{DD UMa} is a $\\delta$ Scuti star whose membership in the \\astrobj{Ursa Major group} is unclear.  The objective of this study is to confirm the membership of \\astrobj{DD UMa} in the \\astrobj{Ursa Major group}, as well as perform a detailed spectral analysis of the star.  Since \\astrobj{DD UMa} is a low-amplitude $\\delta$ Scuti star, we performed a frequency analysis.  We determined fundamental parameters, chemical abundances, and derive a mass and age for the star.  For this study we observed \\astrobj{DD UMa} at the Okayama Astrophysical Observatory with the high-resolution spectrograph HIDES, between the $27^{th}$ of February and the $4^{th}$ March, 2009. Additional observations were extracted from the ELODIE archive in order to expand our abundance analysis. Group membership of \\astrobj{DD UMa} was assessed by examining the velocity of the star in Galactic coordinates. Pulsational frequencies were determined by examining line profile variability in the HIDES spectra. Stellar fundamental parameters and chemical abundances were derived by fitting synthetic spectra to both the HIDES and ELODIE observations.  \\astrobj{DD UMa} is found to be a member of the extended stream of the \\astrobj{Ursa Major group}, based on the space motion of the star.  This is supported by the chemical abundances of the star being consistent with those of \\astrobj{Ursa Major group} members.  The star is found to be chemically solar, with $T_{eff}=7450\\pm150$ K and $\\log g=3.98\\pm0.2$.  We found pulsational frequencies of 9.4 c/d and 15.0 c/d.  While these frequencies are insufficient to perform an asteroseismic study, \\astrobj{DD UMa} is a good bright star candidate for future study by the BRITE-constellation.} ",
        "introduction": "\\label{}    The \\astrobj{Ursa Major group}, also known as the \\astrobj{Sirius supercluster}, is the closest supercluster to the Sun. This group was first identified by the observation that five central stars in Ursa Major were moving towards the same point in space \\citep{proctor, huggins}. After this discovery, a number of other stars were noticed to have similar space motions, some close to Ursa Major and some as far as the southern hemisphere (e.g HD 147584). This leads to distinguishing two subgroups in the \\astrobj{Ursa Major group}: the nucleus (consisting of 13 stars including the five central stars in Ursa Major), and the extended stream. The subgroups are considered to share a common origin, but the members are considered gravitationally unbound. Stars located in the extended stream, as well as in the nucleus, provide information about the formation and evolution of the Ursa Major moving group. \\astrobj{DD UMa} (18 UMa, $m_V$ = $4.^m832$) is a $\\delta$ Scuti star located in the Ursa Major constellation, for which membership in the \\astrobj{Ursa Major group} has not been well defined due to the lack of spectroscopic and photometric observations.  Spectroscopic variability in \\astrobj{DD UMa} was first noticed by \\citet{schesinger}, based on observations from the Allengheny Observatory, who suggested the star was a spectroscopic binary candidate.  \\citet{abt} obtained radial velocity measurements during 13 nights spread between December 1959 and April 1961, and reported that there was no variability in the radial velocity curve of \\astrobj{DD UMa}, thus they concluded \\astrobj{DD UMa} was not a spectroscopic binary.  \\citet{percy} used  \\astrobj{DD UMa} as a comparison star while photometrically observing HR 3775. In these observations they noticed that \\astrobj{DD UMa} showed $\\delta$ Scuti variability, with a $\\sim$0.03 magnitude amplitude and a period of $\\sim$3 hours.  \\citet{horan} observed \\astrobj{DD UMa} for two nights and confirmed the period and amplitude reported by \\citet{percy}. The period derived by \\citet{percy} is the only value available in the literature, and was recently included in work by \\citet{rodriguez}, \\citet{samus}, and \\citet{watson}.   With this paper we supply a modern spectroscopic analysis of \\astrobj{DD UMa}, which is notably absent in the literature. A detailed abundance analysis and frequency analysis of the star is important groundwork for further studies of the star, particularly asteroseismology, and helps place the star in a proper evolutionary context.  A careful assessment of \\astrobj{DD UMa}'s membership in the Ursa Major moving group is also important. An older pulsating star such as \\astrobj{DD UMa} could provide valuable information about the formation and evolution of \\astrobj{Ursa Major group}.  Additionally, \\astrobj{DD UMa} is a candidate target for the BRITE-Constellation\\footnote{http://www.brite-constellation.at/}, which will perform high-precision photometry of bright stars in order to analyse their pulsational properties.  Thus a detailed ground based study of \\astrobj{DD UMa} is valuable in advance of these potential space based observations. ",
        "conclusions": "We conclude that kinematically \\astrobj{DD UMa} is located in the extended stream of the \\astrobj{Ursa Major group}, and that it satisfies the spectroscopic criteria given by \\citet{ammler} for a group member.  \\astrobj{DD UMa} has solar chemical abundances, an effective temperature of $7450\\pm150$ K  and a $\\log g$ of $3.98\\pm0.2$.  From the star's H-R diagram position we infer a mass of $1.72 \\pm0.02$ M$_\\odot$ and an age of $1.05^{+0.10}_{-0.15} \\times10^9$ years, which makes it coeval with the \\astrobj{Ursa Major group}.  As a $\\delta$ Scuti pulsator, we identified a main pulsational frequency of 9.4 c/d, and a second frequency of 15.0 c/d.  \\astrobj{DD UMa} is an unbound and relatively old (compared to the group average) member of the \\astrobj{Ursa Major group}. It will provide important information for subsequent kinematic evolution and origin studies of the \\astrobj{Ursa Major group}. Furthermore, this study provides valuable information for future asteroseismic studies with the BRITE-constellation, which is planned to launch this year. Photometric measurements with the BRITE-constellation will reveal more pulsation frequencies for mode identification, which can be compared to detailed asteroseismic models.       \\bigskip {\\bf \\noindent Acknowledgements} This paper is based on observations using the HIDES spectrograph at the Okayama Astrophysical Observatory (Japan). A. Elmasl\\i\\, gratefully acknowledges the National Astronomical Observatory of Japan (NAOJ) for providing financial support for the observing run (Project no: 09A-12). She also acknowledges Ankara University for providing her a travel grant to the UK and the Open University (UK) for funding accommodation during the analysis of the data."
    },
    "1107/1107.0969_arXiv.txt": {
        "abstract": "{Recent studies have shown that star formation in mergers does not seem to follow the same Schmidt--Kennicutt relation as in spiral disks, presenting a higher star formation rate (SFR) for a given gas column density.} {In this paper we study why and how different models of star formation arise. To do so we examine the process of star formation in the interacting system Arp~158 and its tidal debris.} {We perform an analysis of the properties of specific regions of interest in Arp~158 using observations tracing the atomic and the molecular gas, star formation, the stellar populations as well as optical spectroscopy to determine their exact nature and their metallicity. We also fit their spectral energy distribution with an evolutionary synthesis code. Finally, we compare star formation in these objects to star formation in the disks of spiral galaxies and mergers.} {Abundant molecular gas is found throughout the system and the tidal tails appear to have many young stars compared to their old stellar content. One of the nuclei is dominated by a starburst whereas the other is an active nucleus. We estimate the SFR throughout the systems using various tracers and find that most regions follow closely the Schmidt--Kennicutt relation seen in spiral galaxies with the exception of the nuclear starburst and the tip of one of the tails. We examine whether this diversity is due to uncertainties in the manner the SFR is determined or whether the conditions in the nuclear starburst region are such that it does not follow the same Schmidt--Kennicutt law as other regions.} {Observations of the interacting system Arp~158 provide the first evidence in a resolved fashion that different star--forming regions in a merger may be following different Schmidt--Kennicutt laws. This suggests that the physics of the interstellar medium at a scale no larger than 1~kpc, the size of the largest gravitational instabilities and the injection scale of turbulence, determines the origin of these laws.} ",
        "introduction": "How atomic gas turns molecular and forms stars is one of the questions that is paramount to understanding the process of star formation in galaxies. This process plays a fundamental role in the transformation of baryonic matter and is the main driver of galaxy formation and evolution. It is well known that star--formation depends on the local physical conditions of the gas such as the pressure. It depends on the molecular gas column density by way of the Schmidt--Kennicutt law \\citep{schmidt1959a,kennicutt1998b,kennicutt2007a,leroy2008a,bigiel2008a}. On the large scale it is, however, still not well understood how star formation depends on feedback, spiral density waves or shear, and more generally on the global environment such as in the case of mergers \\citep{barnes2004a,chien2010a,teyssier2010a,bournaud2011a}.  One way to test for the effect of the environment is to use star--forming regions located in collisional debris which are the offsprings of galaxy interactions. Indeed, depending on the parameters of the collision (relative velocity, impact factor, prograde versus retrograde, etc.), a varying amount of gas and stars can be stripped from the parent galaxies and injected into the intergalactic medium. In addition to atomic gas pulled from the parent galaxies, these collisional debris actually contain surprisingly large amounts of molecular gas formed in--situ \\citep{braine2000a,braine2001a,lisenfeld2002a,lisenfeld2004a,petitpas2005a,duc2007a}. As the gas subsequently collapses, star--forming regions are created with masses ranging from a few hundred solar masses, creating OB associations, the ``emission line dots'' \\citep{gerhard2002a,yoshida2002a,sakai2002a,weilbacher2003a,ryan2004a,mendes2004a,cortese2006a,werk2008a,werk2010a} to objects as massive as dwarf galaxies named Tidal Dwarf Galaxies \\citep[TDGs,][]{duc1995a,duc1998a,duc2000a,duc2007a,hancock2007a,hancock2009a}. These objects, even though formed from material once pertaining to their parent galaxies, have a radically different and simpler environment. Recently, \\cite{boquien2007a,boquien2009a} showed that star--formation tracers such as ultraviolet, mid--infrared, and H$\\alpha$, are as reliable in intergalactic star--forming regions as they are in spiral galaxies. \\cite{braine2001a} showed that the depletion timescale of the molecular gas is similar to that of spiral galaxies even though the collision debris they studied have the luminosity, mass and colour of dwarf galaxies. As collision debris probably have a conversion factor close to that of spiral galaxies with a solar neighbourhood metallicity, this eliminates a major source of uncertainty.  There are still many open questions regarding star--formation in collision debris and how the process compares to star formation in spiral disks or in other environments. It has been observed that star formation is particularly efficient in ultraluminous infrared galaxies (ULIRGs) deduced from a high L(TIR)/L(CO) \\citep{solomon2005a}, as well as in galaxies that have a subsolar metallicity \\citep{leroy2006a,gardan2007a,gratier2010a,braine2010b}. Conversely, other systems like intergalactic collisional gas bridges contain molecular gas but form stars inefficiently, such as the UGC 813/6 pair \\citep{braine2004b}. Recent results by \\cite{daddi2010a} and \\cite{genzel2010a} showed that starbursts and mergers follow different modes of star formation compared to more quiescent spiral galaxies. Starbursts appear to be forming stars at a higher rate than spiral galaxies for the same gas column density. However, the discrepancy between the 2 regimes  is increased by the use of a different X$_\\mathrm{CO}$ conversion factor between star--forming galaxies and starbursting ones. Can we observe such a difference within a single interacting system, hence retrieving this result in a resolved fashion? Are there deviations in the Schmidt--Kennicutt law between intergalactic star formation regions and what is seen in galactic disks \\citep{bigiel2008a}?  To address these questions we compare star formation in different parts of a single interacting system, Arp~158, and its tidal debris. This system is particularly suited for this study as it presents evidence of extended star formation. Using an homogeneous dataset on a single system means there are no internal calibration or methodological differences for the study of star formation in the interacting galaxies and their tidal debris. Arp~158 is an intermediate--stage merger in the \\cite{toomre1977a} sequence, with the disks progressively merging but still with 2 clearly separated nuclei. With a recession velocity $cz=4758$~km~s$^{-1}$, Arp~158 is located at a luminosity distance of 62.1~Mpc\\footnote{Value obtained from NED, the NASA Extragalactic Database operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.}, assuming $\\mathrm{H_0}=73$~km~s$^{-1}$~Mpc~$^{-1}$, $\\Omega_\\mathrm{m}=0.27$ and $\\Omega_\\mathrm{\\Lambda}=0.73$, which corresponds to a scale of 292~pc/arcsec. To carry out this work we use multi--wavelength data, ranging from the far--ultraviolet to the radio, to trace star formation, the stellar populations and the atomic gas, also encompassing observations of the molecular component. Optically, the closely interacting galaxies have a boxy shape. They present a bar with 3 nearly aligned bright clumps. The nature of one of them, galactic nuclei or a foreground star, is debated in the literature \\citep{chincarini1973a,dahari1985a}. Prominent dust lanes are visible North of the bar. Two tidal tails can be seen, containing blue concentrations at their tips. The longer one, oriented towards the East--South--East, is a prime Tidal Dwarf Galaxy candidate. The disks of the galaxies seem to be rather edge--on but their morphology is strongly disturbed making a precise assessment difficult. Radio 21~cm observations have shown that the entire system contains $6.5\\times10^9$~M$_\\odot$ of HI, with column density peaks at the tip of one of the tidal arms and towards the western nucleus \\citep{iyer2004a}, providing a large reservoir of gas to fuel star formation throughout the system which yields an infrared luminosity of $\\mathrm{3.1\\times10^{10}~L_\\odot}$ (Sec.~\\ref{ssec:SFR}).  In Sec.~\\ref{sec:observations} we present the multi--wavelength observations, followed by the results in Sec.~\\ref{sec:results} which we discuss in Sec.~\\ref{sec:discussion} before concluding in Sec.~\\ref{sec:conclusion}.  \\section {Observations and data reduction} \\label{sec:observations}  In this section we present the set of multi--wavelength data we are using to carry out this study. For an easier description of the observations we have labelled different regions of interest in the system. These regions have been selected from the inspection of multi--wavelength data and concern the most interesting morphological features. The main characteristics of each object are briefly described in Table \\ref{tab:regions} and they are labelled on images of the system in Fig.~\\ref{fig:system-image}.  \\begin{table*}[!htbp] \\caption{Regions of interest\\label{tab:regions}} \\begin{tabular}{lccl} \\hline Label&$\\alpha$&$\\delta$&Comment\\\\\\hline NE&01:25:22.334&+34:01:32.39&Eastern nucleus\\\\ NC&01:25:20.734&+34:01:29.84&Central nucleus\\\\ NW&01:25:19.467&+34:01:30.08&Western object: star or nucleus\\\\ C1&01:25:20.447&+34:01:25.35&Optical faint blue region at the base of the southern arm\\\\ C2&01:25:16.830&+34:01:37.13&Several compact blue regions to the West, in the northern arm\\\\ C3&01:25:18.670&+34:01:36.30&Several compact blue regions in the northern arm\\\\ C4&01:25:21.893&+34:01:21.54&Faint compact blue regions in the southern arm\\\\ C5&01:25:25.644&+34:01:07.12&Clumpy compact blue regions in the TDG candidate to the East\\\\ \\hline \\end{tabular} \\end{table*}  \\subsection{CO} \\label{ssec:CO}  To trace the molecular gas content of the system we have carried out CO observations with the 30 metre millimetre--wave telescope on Pico Veleta (Spain) run by the Institut de Radio--Astronomie Millim\\'etrique (IRAM). In November 2006 we observed positions C5 and NC with the A and B receivers. The observations of the positions NC, NE, C2 and C3  were done in January 2010 with the EMIR receiver. The CO(1--0) and CO(2--1) transitions at 115 GHz and 230 GHz, respectively, were observed simultaneously and in both polarisations. The redshifted frequencies were 113.4547 and 226.905 GHz, corresponding to a central velocity of $cz = 4800$~km~s$^{-1}$. For the November 2006 observations, we used the 512 $\\times$ 1 MHz filterbanks at 3~mm, one for each polarisation, and the two 256 $\\times$ 4 MHz filterbanks for the two polarisations at 1 mm, yielding instantaneous bandwidths of 1300~km~s$^{-1}$ in all transitions. The observations in January 2010 used the autocorrelator Wilma with a resolution of 2 MHz and for both the 90 GHz and 230 GHz band. The bandwidth was over 10000~km~s$^{-1}$ for the CO(1--0) line and over 5000~km~s$^{-1}$ for the CO(2--1) line. All the  observations were carried out in good weather conditions, the mean system temperatures being for the CO(1--0) transition 190 K (November 2006), and 245 K (January 2010), and for the CO(2--1) transition 330 K (November 2006), and 240 K (January 2010) on the T$_A ^*$ scale. The main beam efficiencies at Pico Veleta were taken to be at 115 GHz (0.74 and 0.77 for the A+B receivers and EMIR respectively), and at 230 GHz (0.54 and 0.58 for the A+B receivers and EMIR respectively). The halfpower beamwidths are about $\\sim22$ and $\\sim11$ arcseconds, corresponding to physical sizes 6.4~kpc and 3.2~kpc. All observations were done in wobbler--switching mode, with a throw in azimuth between 150 and 220 arcseconds. Pointing was checked on a nearby quasar roughly every 90 minutes.  Data reduction was straightforward: the spectra for each position were averaged. In most cases only zero--order baselines (i.e. continuum levels) were subtracted to obtain the final spectra, just in a few spectra a linear baseline had to be subtracted. Position NC was observed during both observing runs and allowed us to check the relative calibration. We found that the velocity integrated intensities were within 10\\% for CO(1--0) and CO(2--1). This constitutes the main source of uncertainty for the 3 brightest regions in CO. The characteristics of the lines are listed in Table~\\ref{tab:CO}.  \\begin{table*}[!htbp] \\caption{Characteristics of the CO lines for the 5 pointings.\\label{tab:CO}} \\begin{tabular}{llllllllll} \\hline Region&I$_\\mathrm{CO(1-0)}$&v$_\\mathrm{CO(1-0)}$&$\\Delta$v$_\\mathrm{CO(1-0)}$&I$_\\mathrm{CO(2-1)}$&v$_\\mathrm{CO(2-1)}$&$\\Delta$v$_\\mathrm{CO(2-1)}$\\\\ &(K~km~s$^{-1}$)&(km~s$^{-1}$)&(km~s$^{-1}$)&(K~km~s$^{-1}$)&(km~s$^{-1}$)&(km~s$^{-1}$)\\\\\\hline NC&$10.08\\pm0.43$&4791&429&$9.0\\pm0.55$&4825&277\\\\ NE&$5.87\\pm0.38$&4900&248&$7.9\\pm0.51$&4911&238\\\\ C2&$0.38\\pm0.09$&4523&146&$<0.62$&4523&146\\\\ C3&$3.60\\pm0.20$&4638&225&$3.07\\pm0.13$&4589&128\\\\ C5&$0.44\\pm0.05$&4955&100&$0.61\\pm0.07$&4920&105\\\\\\hline \\end{tabular} \\tablefoot{The velocity integrated intensities are in K~km~s$^{-1}$ in T$_\\mathrm{mb}$ scale and velocities are in km~s$^{-1}$. The velocity width is taken from the 0--level of the spectra. The errors are 1--$\\sigma$ errors derived from the rms noise of the spectra. A 10\\% calibration error needs to be considered in addition.} \\end{table*}  \\subsection{HI} \\label{ssec:HI} The HI data cube was obtained from the literature \\citep{iyer2004a}. Observations of the atomic hydrogen were performed at the VLA (Very Large Array) in D configuration between 1999-05-14 and 1999-05-17 for a total of 5~h. The velocity resolution is 12.3~km~s$^{-1}$ and the final map resolution is 43.3\\arcsec$\\times$42.2\\arcsec\\ corresponding to a physical size of 12.3~kpc$\\times$12.6~kpc. Data reduction is described in detail in \\cite{iyer2004a}.  \\subsection{Ultraviolet, optical and infrared imaging}  Ultraviolet data have been obtained from the public archives of the GALEX space observatory. Observations were carried out on 2005--11--17 in the FUV (far ultraviolet; $\\lambda_{\\textrm{eff}}$=151~nm) and NUV (near ultraviolet; $\\lambda_{\\textrm{eff}}$=227~nm) bands in the context of the Guest Investigator programme GI2-019 (PI: Koratkar). Arp~158 is located $\\sim3$\\arcmin\\ from the edge of the 1.24$^\\circ$ field-of-view which is centered on NGC~507. The exposure time in FUV (NUV) is 3230~s (3398~s) leading to a sensitivity of 38~nJy~px$^{-1}$ (resp. 26~nJy~px$^{-1}$). The point spread function has a FWHM (full width half maximum) size of $\\sim5$\\arcsec\\ corresponding to a physical size 1.5~kpc.  Shallow optical images have been obtained from SDSS (Sloan Digital Sky Survey) release 7 \\citep{abazajian2009a}, tile 4829 in $u'$, $g'$, $r'$, $i'$ and $z'$ bands. Observations were carried out on 2004--09--15. Standard SDSS calibrations have been applied.  Shallow near--infrared images of the system were obtained from the 2MASS \\citep[2 Micron All Sky Survey,][]{skrutskie2006a} archives. They are used exclusively for visual inspection.  Finally, Arp~158 was observed by our team in the 24~$\\mu$m band on 2009--02--07 using the Spitzer Space Observatory, programme ID 50191. The final exposure time per pixel is 10~s. The frames were processed by the standard Spitzer pipeline version 18.5.0. The processed image presents a smooth gradient in the background which was eliminated subtracting on each row the median of the sky.  Images of the system combining various wavelengths are presented in Fig.~\\ref{fig:system-image} along with the regions of interest identified.  \\begin{figure*}[!http] \\center \\includegraphics[width=\\textwidth]{system_white.pdf} \\caption{Top: optical image in $g'$ (blue channel), $r'$ (green channel) and $i'$ (red channel) bands. The regions of interest listed in Table \\ref{tab:regions} are indicated with a red label. The white dashed rectangles represent the two long slits used for optical spectroscopy. Bottom: FUV (blue channel), NUV (green channel) and 24~$\\mu$m (red channel). The contours in cyan trace the HI column density. The yellow dashed circles represent the CO(1--0) beams and the area in which fluxes have been measured. Due to the large CO beams, not all regions of interest can have their CO emission detected separately.\\label{fig:system-image}} \\end{figure*}  Fluxes were measured in all images within circular apertures matching the CO(1--0) beam (about 22\\arcsec\\ for the half--power beamwidth, Sect. \\ref{ssec:CO} and Fig.~\\ref{fig:system-image}) using the {\\sc phot} procedure from {\\sc iraf}. The background was taken as the mean of the median pixel value in several positions around the interacting system. The same background was subtracted for all apertures. Recent studies of nearby galaxies have shown the importance of background subtraction to derive the star--formation laws \\citep{blanc2009a,rahman2011a,liu2011a}. However, this is most important for regions in galaxies forming stars at a low rate that can be resolved only in nearby galaxies. The distance of Arp~158, 62.1~Mpc, and the size of the regions selected limits the bias introduced by the background subtraction. The uncertainty on the background level is taken as the quadratic mean of the standard deviation of the background level measurements and the mean of the standard deviation of pixels within apertures in which the background is measured. Foreground extinction from the Milky Way was corrected in the UV and in the optical using the \\cite{cardelli1989a} extinction law assuming $\\mathrm{E(B-V)}=0.054$ from NED. The foreground  extinction--corrected fluxes from FUV to 24~$\\mu$m are presented in Table~\\ref{tab:fluxes}.  \\begin{table*}[!htbp] \\caption{UV, optical and IR fluxes in the 5 CO pointing beams.\\label{tab:fluxes}} \\begin{tabular}{ccccccccc} \\hline Region & FUV & NUV & $u$ & $g$ & $r$ & $i$ & $z$ & 24~$\\mu$m\\\\ &($\\mu$Jy)&($\\mu$Jy)&(mJy)&(mJy)&(mJy)&(mJy)&(mJy)&(mJy)\\\\ \\hline NC&$76\\pm4.1$&$189\\pm9.9$&$1.13\\pm0.03$&$4.34\\pm0.09$&$8.76\\pm0.18$&$13.1\\pm0.27$&$17.7\\pm0.39$&$57.3\\pm2.3$\\\\ NE&$101\\pm5.3$&$270\\pm13.8$&$1.51\\pm0.03$&$5.54\\pm0.11$&$10.3\\pm0.21$&$14.0\\pm0.28$&$17.5\\pm0.38$&$139.0\\pm5.6$\\\\ C2&$52\\pm3.0$&$94\\pm5.6$&$0.29\\pm0.01$&$0.82\\pm0.02$&$1.24\\pm0.03$&$1.56\\pm0.05$&$1.72\\pm0.16$&$6.8\\pm0.5$\\\\ C3&$132\\pm6.8$&$245\\pm12.6$&$1.58\\pm0.03$&$4.88\\pm0.10$&$7.46\\pm0.15$&$9.18\\pm0.19$&$10.3\\pm0.26$&$18.6\\pm0.8$\\\\ C5&$31\\pm2.1$&$55\\pm4.1$&$0.17\\pm0.01$&$0.46\\pm0.02$&$0.71\\pm0.02$&$0.93\\pm0.05$&$1.0\\pm0.16$&$2.1\\pm0.4$\\\\ \\hline \\end{tabular} \\tablefoot{The 1--$\\sigma$ error bars take into account uncertainties on the absolute flux calibration and on the background level. Note that the errors are generally dominated by the uncertainty on the flux calibration.} \\end{table*}  \\subsection{Optical spectroscopy}  To gain insight into the physical nature of the selected regions and estimate their metallicity, we have also performed optical spectroscopy. Observations were carried out in service mode under good weather conditions at the William Herschel Telescope. Two long--slit spectra were obtained by our team using the ISIS (Intermediate dispersion Spectrograph and Imaging System) instrument with an exposure time of 1200~s. Each slit is 4\\arcmin\\ long and 2\\arcsec\\ wide. The first slit was aligned along the 3 bright condensations NE, NC, and NW, and the bar, with a position angle of 175.5$^\\circ$. The second one is aligned along the eastern tidal arm and the FUV bright condensation capturing clumps C1, C2 and C5, with a position angle of 15.5$^\\circ$. Each slit was observed using the R600R grating (dispersion of 0.50~\\AA/pixel) from 612~nm to 820~nm in the red arm of ISIS and the R600B grating (dispersion of 0.44~\\AA/pixel) from 360~nm to 540~nm in the blue arm of the instrument. An image of the system in the optical along with the position of the slits is presented in Fig.~\\ref{fig:system-image}. Data reduction was carried out with the {\\sc longslit} package of {\\sc iraf} in a standard way. Flux calibration was carried out using the spectrophotometric star Feige 15 observed the same night. The emission lines were measured with {\\sc splot} within the {\\sc IRAF} package. We present the data in Fig.~\\ref{fig:spectra}.  \\begin{figure*}[!htbp] \\center \\includegraphics[width=\\columnwidth]{spectrum_blue.pdf} \\includegraphics[width=\\columnwidth]{spectra.pdf} \\caption{Left: blue part of the spectrum of the NE clump. Several lines are clearly detected: $[\\mathrm{OII}]_{372.7}$, H$\\gamma$, $[\\mathrm{OIII}]_{436.3}$, H$\\beta$, $[\\mathrm{OIII}]_{495.9}$, and $[\\mathrm{OIII}]_{500.7}$. Unfortunately the blue spectrum for other clumps is not deep enough to detect these emission lines. Right: spectra of clumps NE, NC, NW, C2, and C5. The spectra have been offset to distinguish them more easily. The spectra of clumps NC, NW, C2, and C5 are multiplied by 5 as they are intrinsically much fainter than the spectrum of NE. Except for NW, several lines are clearly detected: $[\\mathrm{NII}]_{654.8}$, H$\\alpha$, $[\\mathrm{NII}]_{658.4}$, $[\\mathrm{SII}]_{671.7}$, and $[\\mathrm{SII}]_{673.1}$. The spectrum of NW shows a broad H$\\alpha$ line in absorption at almost 0--velocity with a weak H$\\alpha$ emission line at the recession velocity of Arp~158.\\label{fig:spectra}} \\end{figure*}  In Table \\ref{tab:spectro} we list the fluxes of the main emission lines we detected in the blue or red spectra depending on their wavelength.  \\begin{table*}[!htbp] \\caption{Spectroscopic observations. \\label{tab:spectro}} \\begin{tabular}{lccccc} \\hline Line&NE&NC&NW&C2&C5\\\\\\hline $[\\mathrm{OII}]_{3727}$& $20.18\\pm0.97$&-- &-- &-- &--\\\\ H$\\gamma$&               $2.11\\pm0.44$&-- &-- &-- &--\\\\ $[\\mathrm{OIII}]_{4363}$&$1.11\\pm0.15$&-- &-- &-- &--\\\\ H$\\beta$&$7.83\\pm0.37$                &-- &-- &-- &--\\\\ $[\\mathrm{OIII}]_{4959}$&$2.81\\pm0.39$&-- &-- &-- &--\\\\ $[\\mathrm{OIII}]_{5007}$&$7.52\\pm0.55$&-- &-- &-- &--\\\\ $[\\mathrm{OI}]_{6300}$&$0.72\\pm0.12$&$0.13\\pm0.08$&-- &-- &--\\\\ $[\\mathrm{NII}]_{6548}$&$4.79\\pm0.11$&$0.57\\pm0.12$&-- &$0.34\\pm0.02$&$0.12\\pm0.02$\\\\ H$\\alpha$       &$49.8\\pm0.15$       &$1.53\\pm0.12$&$0.59\\pm0.68$&$4.73\\pm0.02$&$1.37\\pm0.02$\\\\ $[\\mathrm{NII}]_{6584}$&$16.7\\pm0.12$&$2.03\\pm0.17$&-- &$0.94\\pm0.03$&$0.39\\pm0.02$\\\\ $[\\mathrm{SII}]_{6717}$&$7.42\\pm0.14$&$0.84\\pm0.19$&-- &$0.64\\pm0.02$&--\\\\ $[\\mathrm{SII}]_{6731}$&$5.89\\pm0.17$&$0.64\\pm0.16$&-- &$0.40\\pm0.03$&--\\\\ \\hline EW(H$\\alpha$)&-23&-2.3&-0.2&-174&-44\\\\ \\hline v&5018&4785&--&4513&4878\\\\ \\hline \\end{tabular} \\tablefoot{Line fluxes are in $\\mathrm{10^{-15}\\ erg\\ s^{-1}\\ cm^{-2}\\ \\AA^{-1}}$, the equivalent widths $\\mathrm{EW(H\\alpha)}$ are in \\AA, the velocities v are in km~s$^{-1}$.} \\end{table*} ",
        "conclusions": "In this paper we have studied how properties of star--forming regions vary across an interacting system. To do so we have combined an extensive set of archival and proprietary data tracing the molecular gas (CO), the atomic gas (HI), star formation (FUV and 24~$\\mu$m), and the stellar populations. The interacting system shows a complex morphology, the disks of the 2 colliding galaxies having already interpenetrated. To ascertain the exact nature of the different regions in the interacting system we have also obtained optical spectra. In particular we have obtained a firm identification of the nuclei of the merging galaxies, which was still under debate. One, to the East, exhibits a starburst, the other one hosts an AGN. A third nucleus, to the West, turns out to be a foreground star. A brief description of the regions of interest in Arp~158 is provided hereafter. \\begin{itemize} \\item {\\it NE}: The emission of the eastern nucleus is peaked and is strong at 24~$\\mu$m, suggesting an AGN or a nuclear starburst. The emission line ratios from the BPT diagram \\citep{baldwin1981a} show that it is most likely the latter. The starburst is probably due to gas losing angular momentum as a result of the interaction, and inflowing towards the centre of the galaxy as a consequence. \\item {\\it NC}: Emission line ratios of the central nucleus hint at an AGN, either an extreme LINER or a Seyfert. This unfortunately compromises any accurate measure of the SFR in this regions. \\item {\\it NW}: Optical spectroscopy clearly shows that this is a foreground star with a strong H$\\alpha$ Balmer line in absorption at zero velocity. A faint H$\\alpha$ emission at the redshift of the system is also observed showing that the star may mask a star--forming region. A close inspection of 2MASS near--infrared images also shows that NW is qualitatively more compact than NC and NE. This confirms what \\cite{chincarini1973a} found and closes the question of the nature of this object. \\item {\\it C1}: A peak is visible at 24~$\\mu$m without counterpart in ultraviolet, hinting that it may be a young, dusty star--forming region that has yet to break the surrounding dust cloud. The CO emission of this clump is encompassed by the CO(1--0) beam for the NC pointing. \\item {\\it C2}: The western tip of the system is peaked at 24~$\\mu$m with a faint counterpart in the ultraviolet. CO observations show that it is only marginally detected in CO(1--0) and not detected in CO(2--1), with a detection threshold at 3--$\\sigma$. \\item {\\it C3}: The region is particularly peaked in FUV and NUV but does not present a similar peak at 24~$\\mu$m. However CO emission is clearly detected. \\item {\\it C4}: A faint 24~$\\mu$m peak can be seen without any counterpart in the ultraviolet showing the possible presence of embedded star--forming regions. The NE CO pointing overlaps partially with this region in CO(1--0). \\item {\\it C5}: The TDG candidate is clearly peaked in both the 24~$\\mu$m and the ultraviolet bands. A faint, diffuse UV emission can be seen in the bridge linking it to the main bodies of the interaction. CO observations show a clear detection both in CO(1--0) and in CO(2--1). \\end{itemize}  Studying the Schmidt--Kennicutt law, most regions in the system follow closely the relation found in spiral galaxies by \\cite{bigiel2008a}, with the exception of the nuclear starburst and the tip of one of the tidal tails which show a significantly larger efficiency. Comparisons with the relations derived by \\cite{daddi2010a} exhibit the presence of 2 star formation modes. The detection of these 2 modes in a resolved way hints that their origin is linked to the physics of the interstellar medium at scales no larger than 1~kpc, the size of the largest gravitational instabilities and the injection scale of turbulence. \\label{sec:conclusion}"
    },
    "1107/1107.5582_arXiv.txt": {
        "abstract": "A basic property of objects, like galaxies and halos that form in cosmological structure formation simulations, is their shape. Here, we critically investigate shape determination methods that are commonly used in the literature. It is found that using an enclosed integration volume and weight factors $r^{-2}$ and $r_\\mathrm{ell}^{-2}$ (elliptical radius) for the contribution of each particle or volume element in the shape tensor leads to biased axis ratios and smoothing of details when calculating the local shape as a function of distance from the center. To determine the local shape of matter distributions as a function of distance for well resolved objects (typically more than $\\mathcal{O}(10^4)$ particles), we advocate a method that (1) uses an ellipsoidal shell (homoeoid) as an integration volume without any weight factors in the shape tensor and (2) removes subhalos. ",
        "introduction": "Typically, the distribution of matter in objects that form in cosmological structure formation simulations is crudely described by spherically averaged density profiles \\cite[e.g.][]{1996ApJ...462..563N,1998ApJ...499L...5M}. But real halos are not spherically symmetric and a natural extension is to describe the iso-density contours as surfaces of ellipsoids. There is a wealth of literature with many different methods that are used to measure the local shape of a mass distribution \\citep[e.g.][]{1983MNRAS.202.1159G,1988ApJ...327..507F,1991ApJ...368..325K,1991ApJ...378..496D,1992ApJ...399..405W,1996MNRAS.281..716C,2002ApJ...574..538J,2004IAUS..220..421S,2004ApJ...611L..73K,2004ApJ...616...27B,2005ApJ...627..647B,2006MNRAS.367.1781A,2007ApJ...671.1135K,2007MNRAS.376..215B,2007MNRAS.377...50H,2008ApJ...681.1076D,2008MNRAS.385.1859W,2010ApJ...720L..62K,2010MNRAS.405.1119K,2011ApJ...734...93L,2011arXiv1104.1566V}. Their common goal is to recover the iso-density surfaces of the underlying matter distribution. Other characteristics, such as the potential, can also be used to describe the objects \\citep[e.g.][]{2004IAUS..220..421S,2007MNRAS.377...50H,2010ApJ...720L..62K,2011ApJ...734...93L}.  Unfortunately, the literature lacks a systematic comparison of the different methods - especially under controlled conditions where the exact shape is known. For some notable exceptions, see e.g. \\cite{2006MNRAS.367.1781A,2004ApJ...611L..73K}. But as far as we know, there is no publication that investigates the different methods under controlled conditions with known shape as it is done in this paper. Presumably, many of the quantitative discrepancies in the literature originate in the various methods that are used for determining the shape. This work is intended to shed some light on the effects and systematics of the various methods that are based on an iterative procedure that uses a shape tensor with different weighting schemes and integration volumes. The influence of the local mass density profile on the capability of the shape finding method to recover the iso-density contours is also investigated. ",
        "conclusions": "} \\tablehead{\\colhead{Method} & \\colhead{$w(\\vec{r})$} & \\colhead{$V$}} \\startdata S1 & 1 & ellipsoidal shell\\\\ S2 & $r^{-2}$ & ellipsoidal shell\\\\ S3 & $r_\\mathrm{ell}^{-2}$ & ellipsoidal shell\\\\ E1 & 1 & enclosed ellipsoid\\\\ E2 & $r^{-2}$ & enclosed ellipsoid\\\\ E3 & $r_\\mathrm{ell}^{-2}$ & enclosed ellipsoid \\enddata \\end{deluxetable}  We concentrate on 6 different methods for determining the shape of a matter distribution (see also Table \\ref{tab:methodsummary}). These methods differ by using a different integration volume $V$ and different weight functions $w(\\vec{r})$. For calculating the \\textit{local} shape at a distance $r_\\mathrm{ell}$, in the methods with a starting letter S, the integration is over an ellipsoidal shell (homoeoid) volume centered at $r_\\mathrm{ell}$ (in logarithmic space). In the methods with first letter E, the integration is over the whole enclosed ellipsoidal volume within $r_\\mathrm{ell}$. For the different weight functions $w(\\vec{r})$, we use (1) $w(\\vec{r}) = 1$, (2) $w(\\vec{r}) = r^{-2}$ and (3) $w(\\vec{r}) = r_\\mathrm{ell}^{-2}$. The elliptical radius is given by \\begin{equation}\\label{eq:rell} r_\\mathrm{ell} = \\sqrt{x_\\mathrm{ell}^2 + \\frac{y_\\mathrm{ell}^2}{(b/a)^2} + \\frac{z_\\mathrm{ell}^2}{(c/a)^2}} \\end{equation} where $(x_\\mathrm{ell},y_\\mathrm{ell},z_\\mathrm{ell})$ are the coordinates of the volume element or particle in the eigenvector coordinate system of the ellipsoid, i.e. $r_\\mathrm{ell}$ corresponds to the semi-major axis $a$ of the ellipsoid surface through that particle or volume element. Additionally, we also check for the importance of the removal of subhalos. Cases where we removed the subhalos are marked with a --, cases where they remained by a +.  In order to calculate the local shape at a distance $r_\\mathrm{ell}$ from the center, we use an iteration method \\citep[e.g.][]{1991ApJ...368..325K,1991ApJ...378..496D,1992ApJ...399..405W} and start with a spherically symmetric integration volume (shell or sphere). Then the shape tensor is calculated according to the different methods. By diagonalizing $\\mat{S}$ we get the eigenvectors and eigenvalues at distance $r_\\mathrm{ell}$. The eigenvectors give the directions of the semi-principal axes. The eigenvalues of $\\mat{S}$ for the method S1 are $a^2/3$, $b^2/3$ and $c^2/3$ where $a$, $b$ and $c$ are the semi-principal axes with $a \\geq b \\geq c$ -- at least in the thin homoeoid approximation where the density is uniform. Hence, the square roots of the eigenvalues are proportional to the lengths of the semi-principal axes for method S1 and we can readily calculate the axis ratios $b/a$ and $c/a$. For method S3 we expect to get the same axis ratios as for method S1 since dividing by the semi-major axis $a=r_\\mathrm{ell}$ squared, which is a constant for a thin ellipsoidal shell, just changes the geometrical meaning and normalization of the eigenvalues but not the axis ratios.  For the other methods it is not clear what the detailed geometrical meaning of the eigenvalues is. For the methods that use the enclosed ellipsoidal volume, this will also depend on the mass density profile. The $r^{-2}$ weighting projects the volume elements d$V$ onto the unit sphere. This projection complicates the physical interpretation of this method. It is generally assumed though that the eigenvalues of $\\mat{S}$ in these cases are still proportional to the semi-major axes squared. Hence, we calculate the axis ratio for the other methods the same way as for methods S1 and S3 -- as it is generally done in the literature.  We then keep the length of the semi-major axis fixed (but the orientation can change) and calculate $\\mat{S}$ again by summing over all particles within the new deformed integration volume (homoeoid or ellipsoid) with semi-major axis $a=r_\\mathrm{ell}$ and axis ratios $b/a$ and $c/a$ but with the new orientation. For the shape determination we allow volume elements or particles to be in several bins/shells. Of course this is naturally the case when using an enclosed ellipsoidal volume. It is also necessary when using an ellipsoidal shell since neighboring shells can overlap due to slightly different orientation and axis ratios. This iteration is repeated until convergence is reached. As a convergence criterion we require that the fractional difference between two iteration steps in both axis ratios is smaller than $10^{-3}$.  For methods using the shape tensor, it is important to use an iteration method that allows the algorithm to adapt the integration volume to the a priori unknown shape of the object. Often one also finds in the literature that no iteration procedure is used and just a simple spherical shell or enclosed sphere is used as the integration volume in order to calculate the shape \\citep[e.g.][]{1983MNRAS.202.1159G,1988ApJ...327..507F,1996MNRAS.281..716C,2004ApJ...616...27B,2005ApJ...627..647B,2010MNRAS.405.1119K}. To us the physical meaning of the outcome of such a procedure is unclear and we do not further pursue it here.  A further method for calculating the shape of contours is by selecting particles by their local density \\citep[e.g.][]{2002ApJ...574..538J,2008MNRAS.385.1859W,2011arXiv1104.1566V} or potential \\citep[e.g.][]{2004IAUS..220..421S,2007MNRAS.377...50H,2010ApJ...720L..62K}. There, no iteration procedure is needed.  Often one also finds in the literature, that the moment of inertia tensor $\\mat{I}$ (Equation (\\ref{eq:mit})) in combination with an enclosed ellipsoidal integration volume is used for calculating the axis ratios. This procedure assumes relations between the eigenvalues and semi-principle axes that are strictly valid only for a uniform ellipsoid or homoeoid \\cite[e.g.][]{2007MNRAS.376..215B,2008MNRAS.385.1859W}. For a thin homoeoid this is fine (under that assumption that the local density is constant in the shell) but for the enclosed ellipsoidal integration volume, the result is made equivalent to the method that just uses the shape tensor by construction.    A widespread method used in the literature is method E3 \\citep[e.g.][]{1991ApJ...378..496D,2006MNRAS.367.1781A,2007ApJ...671.1135K}. By using an enclosed integration volume, this method picks up information from the inner regions that can have different shapes and orientation. If one is interested in the local shape, then we find that method S1 is clearly a better choice than method E3.  Method E1 \\citep[e.g.][]{1991ApJ...368..325K} is doing relatively well compared to its differential version S1. This is due to the fact that the contribution in the shape tensor (Equation (\\ref{eq:shapetensor})) is dominated by particles or volume elements with the largest distance from the center. This method also shows systematic shifts (see for example Figure \\ref{fig:shapemedian}) and smoothing when compared to method S1. Therefore, the differential version S1 should be preferred over the E1 method that uses the enclosed ellipsoidal volume.  Unfortunately, methods S1 or S3 are not yet in widespread use in the literature. \\cite{2004ApJ...611L..73K} used method S3 and also found that using the enclosed volume is sensitive to the distribution of particles in the enclosed region. Unfortunately, they did not present the details of the tests in their work. \\cite{2008ApJ...681.1076D} and \\cite{2011ApJ...734...93L} are also advocating method S1. While \\cite{2008ApJ...681.1076D} do not further motivate their choice, \\cite{2011ApJ...734...93L} found from visual comparison that using a differential method in 2 dimensions gives reliable ellipsoidal fits to X-ray isophotes."
    },
    "1107/1107.5061_arXiv.txt": {
        "abstract": "We report Submillimeter Array (SMA) observations of CO (J=2--1, 3--2 and 6--5) and its isotopologues ($^{13}$CO J=2--1, C$^{18}$O J=2--1 and C$^{17}$O J=3--2) in the disk around the Herbig~Ae star HD~163296 at $\\sim2''$ (250 AU) resolution, and interpret these data in the framework of a model that constrains the radial and vertical location of the line emission regions. First, we develop a physically self-consistent accretion disk model with an exponentially tapered edge that matches the spectral energy distribution and spatially resolved millimeter dust continuum emission. Then, we refine the vertical structure of the model using wide range of excitation conditions sampled by the CO lines, in particular the rarely observed J=6--5 transition. By fitting $^{13}$CO data in this structure, we further constrain the vertical distribution of CO to lie between a lower boundary below which CO freezes out onto dust grains (T $\\lesssim 19$ K) and an upper boundary above which CO can be photodissociated (the hydrogen column density from the disk surface is $\\lesssim10^{21}$ cm$^{-2}$). The freeze-out at 19 K leads to a significant drop in the gas-phase CO column density beyond a radius of $\\sim$155~AU,  a ``CO snow line'' that we directly resolve. By fitting the abundances of all CO isotopologues, we derive isotopic ratios of $^{12}$C/$^{13}$C, $^{16}$O/$^{18}$O and $^{18}$O/$^{17}$O that are consistent with quiescent interstellar gas-phase values. This detailed model of the HD~163296 disk demonstrates the potential of a staged, parametric technique for constructing unified gas and dust structure models and constraining the distribution of molecular abundances using resolved multi-transition, multi-isotope observations. ",
        "introduction": "The disks around pre-main sequence stars are the reservoirs of raw material that represent the initial conditions for the formation of planetary systems. The spatial distribution of mass in these disks is a fundamental property, as it sets a constraint on the contents and orbital architectures of planetary systems.  In that sense, measurements of disk densities and temperatures can provide strong, albeit complex, constraints on planet-forming scenarios. Because the dominant constituent of these disks is believed to be cold H$_2$ gas that is generally unobservable, our knowledge of many disk properties (including mass) relies on the observation and interpretation of minor constituents. Dust has naturally been the focus of the most work on disk structure and evolution, as it dominates the opacity and emits bright continuum  radiation \\citep[e.g.,][]{andrews09,andrews10}. However, converting those continuum surface brightnesses to disk (gas) densities is challenging due to the large uncertainties in the assumed dust properties (composition, size distribution, etc.) and the dust abundance relative to the gas, which may deviate substantially from the canonical interstellar value (0.01) and vary spatially in the disk. In principle, spectral line emission from trace gas species can provide complementary constraints on the disk structure, but the interpretation of the observations can also be difficult.  These emission lines depend on the gas abundances that are set by chemical and physical processes, as well as the excitation conditions set by the local densities, temperatures, and incident radiation fields.  With a dataset of sufficient quality, these complexities can be leveraged into powerful probes of otherwise inaccessible disk characteristics \\citep[e.g.,][]{kamp_t10}.  The chemical property of CO, the most abundant molecule after H$_2$, is thought to be well understood in disks. The spatial distribution of gas-phase CO is expected to be shaped by photodissociation both near the star and high in the disk atmosphere, as well as by a depletion process when it freezes onto dust grains at low temperatures (below $\\sim$20~K) at large radii and deep in the disk interior \\citep[e.g.,][]{aikawa_m96,aikawa_n06,gorti_h08}. The peak abundance of CO is also predicted to be almost independent of radius, around 10$^{-4}$, the typical value in molecular clouds (\\citealp{aikawa_n06}). However, observational studies of CO emission from disks have yielded puzzling results. In general, the CO in disks is found to be under-abundant by factors of 10--100 compared to molecular clouds, a deficit usually attributed to CO depletion in cold, dense gas (\\citealp{dutrey_g94, dutrey_g96, dutrey_g97, vanzadelhoff_v01, thi_v04}). However, multi-transition and multi-isotope CO studies report a wide range of relative depletions of CO and its isotopologues.  For example, \\citet{dartois_d03} finds that all of the CO isotopologues are depleted by an identical uniform factor of $\\sim$10, but \\citet{pietu_d07} finds in a small sample of disks that the $^{12}$CO/$^{13}$CO ratios in the outer part of disks are much lower than the standard $^{12}$C/$^{13}$C ratio in the solar neighborhood: these low values were attributed to significant carbon fractionation. At the same time, both of these studies suggested that the outer radius derived from $^{13}$CO is smaller than the one derived from $^{12}$CO indicating their $^{12}$CO/$^{13}$CO ratios are abnormally large at the outer disk edge of their models, which leads to speculations on the importance of selective photo-dissociation due to self-shielding in the photolysis of gaseous CO (\\citealp{dutrey_g07}). They also find that a large amount of CO remains gaseous at temperatures as low as 10 K. Possible explanations for the presence of the inferred cold CO include vertical mixing (\\citealp{aikawa07}) and photodesorption off dust grains (\\citealp{hersant_w09}). All above depict a very complex picture of the distribution of CO gas in the disks, which remains to be explained by detailed chemical models combined with realistic physical disk structures. However, it is unclear, how much of these variations depend on the individual sources, and how much, if any, depends on the modeling approach.  A robust methodology for interpreting CO observations is a high priority for our efforts to constrain the total gas distribution in disks, and to provide a basis for interpreting observations of more complex molecules.  In this study, we develop a modeling framework that is able to account for observations of both the dust continuum and CO line emission from a protoplanetary disk. Specifically, we combine a previously-established dust emission modeling formalism with multiple transitions, spatially resolved CO line data (CO J=2--1, 3--2 and 6--5, $^{13}$CO J=2--1, C$^{18}$O J=2--1 and C$^{17}$O J=3--2) to construct a simplified and internally self-consistent model of the disk around the Herbig  Ae star HD~163296.  At a distance of $\\sim$122 pc (\\citealp{perryman_l97}), this 2.3 M$_{\\odot}$ star (spectral type A1; age $\\sim$4 Myr) harbors a large disk with strong millimeter continuum and molecular line emission (\\citealp{qi01,thi_v04,natta_t04}). It does not seem to be associated with any known star-formation region, dark clouds, or reflection nebulae. The disk around HD~163296 exhibits a scattered light pattern extending out to a radius of $\\sim$500 AU \\citep{grady_d00}, perpendicular to a bipolar microjet traced by a chain of HH knots visible in coronagraphic observations (\\citealp{devine_g00,grady_d00,wassell_g06}).  Millimeter and submillimeter interferometric observations have explored this large disk and its Keplerian velocity field in low-J CO lines (\\citealp{mannings_s97, isella_t07}). The large disk size and strong molecular emission make this system an excellent target for resolved millimeter observations and detailed modeling.  This paper is organized as follows. We describe the CO and mm continuum data in \\S 2, and present the main observational results in \\S 3. Our extensive modeling effort is detailed in \\S 4. We discuss the modeling results and compare with previous work on this subject in \\S 5, and provide a brief summary in \\S 6. ",
        "conclusions": "\\subsection{Modeling Dust Emission and CO Line Emission}  We have modeled multiple emission lines of CO and its isotopologues from the disk around HD~163296 in the context of an accretion disk model structure, grounded in observations of the broadband SED and resolved millimeter continuum emission.  The goal of this modeling effort is to develop a more consistent examination of the connection between the gas and dust phases in the disk. While our modeling framework does not treat the complexity of a completely self-consistent, simultaneous description of the energy balance and chemistry between the gas and dust, it effectively employs a set of parameters that can retrieve molecular abundance information in a way that captures the essential character of the layered disk structures predicted by those more sophisticated models.  Most importantly, our approach directly addresses two common issues that are noted in much simpler structure models: (1) a radius (or size) discrepancy between the dust and CO emission, and (2) the degeneracy in the vertical temperature structure for models based solely on the dust emission (i.e., the SED).  A longstanding problem with disk models has been the seemingly different radial distributions of dust and gas when each is considered independently (\\citealp{dutrey_g07}). \\citet{isella_t07} presented multi-wavelength millimeter continuum and CO isotopologue observations of the disk around HD~163296 and found a significant discrepancy between the outer radius derived for the dust continuum ($200 \\pm 15$ AU) and that derived from CO emission ($540 \\pm 40$ AU) in their truncated power law models. However, \\citet{hughes_w08} showed that models with tapered outer edges can naturally reconcile the apparent size discrepancy in dust and gas millimeter imaging. The successful fitting of CO isotopologues in the HD~163296 disk using the SED-based disk model with an exponentially tapered outer edge, without invoking an unknown or unconstrained chemical effect, provides new support for the necessity of including this feature in outer disk structure.  The SED alone does not provide any direct information on the temperature structure of the intermediate disk layers where the CO (and other molecular) emission is generated.  In our modeling framework for the HD~163296 disk, the SED (and millimeter continuum images) can be fit equally well for a wide range of vertical temperature/density profiles, highlighting the degeneracy of the dust data with the parameter $z_{\\rm big}$, the height marking the transition between the small grains in the disk atmosphere and the big grains concentrated toward the midplane.  However, we have found that resolved observations of optically thick CO lines at a range of excitations can be used to place stringent constraints on the vertical temperature structure.  Previous analysis of the CO J=2--1 and 3--2 lines from the HD~163296 disk suggested that gas temperatures were always higher than 20~K, ruling out CO freeze-out as a cause for the depletion of CO abundances \\citep{isella_t07}.  But as we discussed in \\S 4.2 (see Figure \\ref{fig:modelspec}), the small excitation leverage between those low-lying transitions is not a strong discriminant of the temperature profile.  Here, we make use of the higher-excitation J=6--5 line to better constrain the vertical structure of the disk, and find a colder midplane that is consistent with significant CO freeze-out.  Although observations of these various CO transitions is expensive, there are likely other molecules that emit at nearby frequencies and can be used to trace a sufficient range of excitation conditions (e.g., HCO$^+$).  The essential point is that the temperature structure in the intermediate layers of a disk can only be constrained well using several optically thick emission lines that probe a range of excitation.   \\subsection{The ``CO Snow Line\"}  Theoretical models of the molecular abundances in protoplanetary disks predict three distinct vertical layers \\citep[see][]{bergin_a07}.  At large heights in the disk atmosphere, temperatures are high and molecular abundances are comparable to those in a standard photon-dominated region (PDR).  At intermediate heights, between the midplane and atmosphere, warm temperatures are suitable for high gas-phase abundances of many molecules.  And at the deepest layers near the midplane, temperatures are low enough that many molecules are depleted from the gas phase and frozen onto the mantles of dust grains.  \\citet{aikawa_n06} and others, have used this kind of layered physical structure in their chemical reaction network models and shown that CO is expected to be abundant in the warm layer, with $f_{\\rm CO} \\sim 10^{-4}$ almost independent of radius.  The boundaries of that layer are defined by the CO freeze-out in the midplane and the photodissociation of CO by high-energy stellar photons in the disk atmosphere.  The modeling of the SMA observations of the HD~163296 disk provides direct support for this basic model structure. Our estimate of the upper boundary $\\sigma_s$ is in excellent agreement with models of the upper PDR layer, where CO is photodissociated into its atomic constituents by stellar UV and/or X-ray radiation at column densities $\\lesssim10^{21}$ cm$^{-2}$ \\footnote{A total gas column of 10$^{21}$ cm$^{-2}$ corresponds to 1 mag of extinction (A$_v=$1) at visible wavelengths when the gas-to-dust ratio is 100. But in our model the CO photodissociation front is located at A$_v\\ll$1 because of the large dust depletion factor as measured with $\\epsilon=0.003$ in the upper disk layers.} \\citep{aikawa_n06,gorti_h08}. The lower boundary that best fits the $^{13}$CO data, corresponding to a temperature of 19~K, is also in excellent agreement with the laboratory studies of CO freeze-out temperature onto dust grains (\\citealp{collings_d03,bisschop_f06}). The main effect of CO photodesorption \\citep{oberg_f07,hersant_w09} and other related non-thermal ice desorption mechanisms is to push the lower boundary to colder temperatures. Since our best-fit CO lower boundary is consistent with thermal desorption, there is no evidence for efficient CO photodesorption in the disk of HD~163296.  Since the term ``snow line\" is often discussed in planetary formation as some point where the temperature in the midplane would drop below the water ice sublimation level, e.g. in the ``minimum-mass solar nebular\" model prescribed by \\citet{hayashi81}, we propose to use the term ``CO snow line\" to indicate where CO freezes out in the disk. It was generally believed that the dust temperature in the disks of Herbig Ae stars is high enough that even the temperature close to the disk midplane is above the temperature of freeze-out of CO (\\citealp{dutrey_g07, jonkheid_d07}). In this work, the primary line of evidence for CO freeze-out in the disk of HD~163296 comes from the detection of sharp reduction of the effective $^{13}$CO J=2--1 emission area in the outer disk. This is demonstrated in Figure~\\ref{fig:zcentco} where the $^{13}$CO visibility is best fit with the lower boundary at 19 K, indicating no or very few $^{13}$CO emission from the midplane area below this boundary, as shown in Figure~\\ref{fig:structure}. The CO snow line marks drop off at least two orders of magnitude in the column density of $^{13}$CO beyond 155 AU, which is clearly resolved by our interferometric observations. However, the emission lines from CO and its isotopologues are still detectable out to large radii, e.g. $>$500 AU, which indicates that there are still substantial amount of CO toward surface layer of the outer disk where the temperature is higher than 19 K.  The presence of this CO snow line has a direct impact on some key elements of chemical networks in a disk, especially related to the ionization fraction. The CO abundance is expected to be linearly correlated with the HCO$^+$ abundance, but has a strong anti-correlation with N$_2$H$^+$ since the main destruction pathway for the latter is in reactions with CO (\\citealp{jorgensen_s04}).  In the midplane, the depletion of CO beyond the snow line should enhance the deuterated molecular ions, H$_2$D$^+$, D$_2$H$^+$ and D$_3^+$, such that they become the most abundant ions in the midplane (\\citealp{ceccarelli_d05}).  So, resolving the CO snow line will eventually help constrain the ionization fraction and the extent of deuterium fractionation in the disk interior.  \\subsection{CO Isotope Ratios}  The model developed here indicates that the gas in the intermediate layer of the HD~163296 disk has CO isotopologue abundance ratios that are consistent to those found in the molecular interstellar medium.  That finding is in contrast with some previous studies of this and other disks that instead suggested an increasing amount of depletion from $^{12}$CO to $^{13}$CO to C$^{18}$O \\citep[e.g.,][]{isella_t07,dutrey_g94,dutrey_g96}.  That difference is likely a manifestation of the underlying modeling approach, as well as the constraining power now available to us from the high-excitation J=6--5 line and the rare isotopologues. When we include CO freeze-out, our models also show no evidence for radial variations of the CO abundance with respect to the H$_2$ densities inferred (indirectly) from the dust emission.  Previous studies have found (much) steeper CO surface density profiles compared to dust (e.g., Pietu et al.~2007, Dutrey et al.~2008).  To compare with those findings, we have used an alternative power-law prescription for the radial abundance fractions, $f_i(r) \\propto r^{p_i}$, and attempted to re-fit our data.  The resulting best-fit values of $p_i$ for $^{13}$CO and C$^{18}$O are $0.0\\pm0.1$ and $0.1\\pm0.2$, respectively: both consistent with a radially constant abundance profile.  So, when taking into account CO freeze-out at low temperatures, there is no clear evidence that the radial distribution of CO deviates from the underlying total gas surface density.  The steep slopes inferred in previous work may be a manifestation of the sharp abundance drop beyond the CO snow line.  This model demonstrates that, given sufficiently strong constraints on the vertical temperature gradient, a unified model of the gas and dust disk can be constructed that conforms to standard ISM-based assumptions about the isotopic abundances and the CO chemistry.  \\citet{visser_v09} present a photodissociation model for CO isotopologues including newly updated depth-dependent and isotope-selective photodissociation rates. They find grain growth in circumstellar disks can enhance $N(^{12}{\\rm CO})/N({\\rm C^{17}O})$ and $N(^{12}{\\rm CO})/N({\\rm C^{18}O})$ ratios by a factor of ten relative to the initial isotopic abundances. We have fit the $\\sigma_s$ and fractional abundances for C$^{17}$O and C$^{18}$O  and found that different pairs of $\\sigma_s$ and f(C$^{18}$O) or f(C$^{17}$O) can fit the data equally well, i.e. we can not constrain $\\sigma_s$  due to the limited signal-to-noise ratios of C$^{18}$O J=2--1 and C$^{17}$O 3--2 data. Future observations with greater sensitivity in the CO isotopologues will be essential to investigate the isotope-selective photodissociation on the distributions of CO isotopologues and hence the local isotopic ratios at different layers of the disks.  \\subsection{Comparison of HD~163296 and TW Hya}  In a series of papers, \\citet{qi_h04,qi_w06} modeled the CO J=2--1, 3--2, and 6--5 line emission from the disk around the cooler young star TW Hya.  There, the best fit model to the optically thick CO J=2--1 and 3--2 lines tended to underestimate the CO 6--5 emission.  To fit all three CO lines simultaneously required additional heating of the disk surface, suggested to be the result of the intense X-ray irradiation field for that source.  No such additional heating is required to explain the HD~163296 data.  On the contrary, the CO J=6--5 emission required a lower disk interior temperature compared to those predicted by the typical accretion disk models based on the SED, forcing us to modify the $z_{\\rm big}$ parameter to fit the data.  There are at least two plausible reasons for the different vertical structures we infer for these disks.  First, is the high-energy radiation field from the two stars. {\\it Chandra} X-ray observations reveal a point-like object within $0\\farcs25$ (30 AU) of HD~163296, with an X-ray luminosity of $4\\times10^{29}$ ergs s$^{-1}$ (\\citealp{swartz_d05}), about 5$\\times$ lower than that from TW Hya (\\citealp{kastner_h99}).  Obviously the suggestion that X-ray heating is dependent on the stellar type will require a much larger sample to confirm, as well as more detailed radiative transfer modeling of the high-energy heating of these disks.  However, it is an exciting prospect that the interior structures of these disks may be very different depending on the connection between the stellar mass and the irradiation environment, with implications both for disk evolution and chemistry.  A second possibility may lie with the level of turbulence in these disks.  In our model, the vertical structure is governed by the settling of large dust grains.  The timescale for gravitational sedimentation is expected to be short, with big grains reaching the midplane in less than $\\sim$1 Myr (\\citealp{dullemond_d04a}) unless the dust is effectively stirred up by turbulence.  The more vertically extended population of big grains required to model the HD~163296 disk might be commensurate with more efficient stirring, or rather inhibited settling.  \\citet{hughes_w11} suggest a 300 m s$^{-1}$ turbulent linewidth in the CO layer of the HD~163296 disk, much larger than the $\\lesssim$40 m s$^{-1}$ they find for the TW Hya disk.  Perhaps turbulent stirring is responsible for lofting the big grains above the midplane in the HD~163296 disk, enhancing the radiative cooling."
    },
    "1107/1107.0644_arXiv.txt": {
        "abstract": "We present ground-based MDM $V$-band and \\textit{Spitzer}/IRAC 3.6$\\mu$m-band photometric observations of the 72 representative galaxies of the \\sauron\\ Survey. Galaxies in our sample probe the elliptical E, lenticular S0 and spiral Sa populations in the nearby Universe, both in field and cluster environments. We perform aperture photometry to derive homogeneous structural quantities. In combination with the \\sauron\\ stellar velocity dispersion measured within an effective radius (\\se), this allows us to explore the location of our galaxies in the colour-magnitude, colour-\\se, Kormendy, Faber-Jackson and Fundamental Plane scaling relations. We investigate the dependence of these relations on our recent kinematical classification of early-type galaxies (i.e. Slow/Fast Rotators) and the stellar populations. Slow Rotator and Fast Rotator E/S0 galaxies do not populate distinct locations in the scaling relations, although Slow Rotators display a smaller intrinsic scatter. We find that Sa galaxies deviate from the colour-magnitude and colour-\\se\\ relations due to the presence of dust, while the E/S0 galaxies define tight relations. Surprisingly, extremely young objects do not display the bluest $(V-[3.6])$ colours in our sample, as is usually the case in optical colours. This can be understood in the context of the large contribution of TP-AGB stars to the infrared, even for young populations, resulting in a very tight $(V-[3.6])-$\\se\\ relation that in turn allows us to define a strong correlation between metallicity and \\se. Many Sa galaxies appear to follow the Fundamental Plane defined by E/S0 galaxies. Galaxies that appear offset from the relations correspond mostly to objects with extremely young populations, with signs of on-going, extended star formation. We correct for this effect in the Fundamental Plane, by replacing luminosity with stellar mass using an estimate of the stellar mass-to-light ratio, so that all galaxies are part of a tight, single relation. The new estimated coefficients are consistent in both photometric bands and suggest that differences in stellar populations account for about half of the observed tilt with respect to the virial prediction. After these corrections, the Slow Rotator family shows almost no intrinsic scatter around the best-fit Fundamental Plane. The use of a velocity dispersion within a small aperture (e.g. \\re/8) in the Fundamental Plane results in an increase of around 15\\% in the intrinsic scatter and an average 10\\% decrease of the tilt away from the virial relation.\\looseness-2 ",
        "introduction": "\\label{sec:intro}  Galaxies are fundamental building blocks of our universe, and our knowledge of their distribution, structure and dynamics is closely tied to our general understanding of structure growth. So-called scaling relations, that is correlations between well-defined and easily measurable galaxy properties, have always been central to our understanding of nearby galaxies. With high redshift studies now routine, scaling relations are more useful than ever, allowing us to probe the evolution of galaxy populations over a large range of lookback times \\citep[e.g.][]{bell04,conselice05,ziegler05,saglia10}.  The colour-magnitude relation (CMR) was already recognised in the sixties and seventies \\citep{devac61,sandage72,vs77} and has served as an important benchmark for theories of galaxy formation and evolution since \\citep[e.g.][]{bower92b,bell04,bernardi05}. The main drivers are thought to be galaxy metallicity, which causes more metal rich galaxies to be redder, and age, causing younger galaxies to be bluer. Galaxies devoid of star formation are thought to populate the red sequence, while star-forming galaxies lie in the blue cloud \\citep[e.g.][]{baldry04}. The dichotomy in the distribution of galaxies in this relation has opened a very productive avenue of research to unravel the epoch of galaxy assembly \\citep[e.g.][]{delucia04,andreon06,arnouts07}.  Since its discovery \\citep{dd87,dressler87}, the Fundamental Plane (FP) has been one of the most studied relations in the literature. Given its tightness, like many other scaling relations the FP was quickly envisaged as a distance estimator as well as a correlation to understand how galaxies form and evolve (e.g.\\ \\citealt{saglia93, jfk96, pahre98, kelson00, bernardi03b, vdwel04, holden05, macarthur09}). It is widely recognised that the FP is a manifestation of the virial theorem for self-gravitating systems averaged over space and time with physical quantities total mass, velocity dispersion, and gravitational radius replaced by the observables mean effective surface brightness (\\mue), effective (half-light) radius (\\re), and stellar velocity dispersion ($\\sigma$). Since velocity dispersion and surface brightness are distance-independent quantities, contrary to effective radius, it is common to express the FP as $\\log($\\re$)\\,=\\,\\alpha\\log(\\sigma)+\\beta$\\mue$+\\gamma$, to separate distance-errors from others. If galaxies were homologous with constant total mass-to-light ratios, the FP would be equivalent to the virial plane and be infinitely thin, with slopes $\\alpha=2$ and $\\beta=0.4$ (in the notation used here). By studying the intrinsic scatter around the FP, one can study how galaxy properties differ within the observed sample.  Some projections of the FP, known earlier in time, have also been widely studied. \\citet{kormendy77} found that the surface brightness density (i.e. mean surface brightness) of a galaxy changes as a function of its size. This relation is usually known as the Kormendy relation (hereafter KR). The correlation is such that larger galaxies have lower surface brightness densities, compared to their smaller counterparts. The size-luminosity relation (SLR) is widely used to establish the size evolution of galaxies as a function of redshift \\citep[e.g.][]{trujillo06,dokkum08}. Finally, the last projection of the FP that we consider in this paper is the Faber-Jackson relation (\\citealt{fj76}; hereafter FJR), which relates the luminosity of a galaxy to its stellar velocity dispersion.  The highest quality and best understood scaling relations in the optical/near-infrared are the ones for early-type galaxies. This is because observationally they are much simpler than spirals, with less complicated star formation histories, and less extinction by dust, and thus tighter scaling relations \\citep[e.g.][]{ls10}. It is for that reason that often the two groups are treated separately in physically similar relations: the prime example being the relations between the central stellar velocity dispersion and absolute luminosity of elliptical galaxies (FJR), and the rotation velocity and absolute luminosity of disc galaxies \\citep{tf77}. In an attempt to unify properties of these two groups, spiral galaxies are usually studied in terms of their bulge and disc properties. The resemblance of bulges to ellipticals has lead to their inclusion in the scaling relations of early-type systems \\citep[e.g.][]{bender92,khos00,fb02}, although they often reveal a much larger scatter and show, on average, an offset with respect to the relations of early-type galaxies. While this might not be surprising due to the mentioned effects of (younger) stellar populations and dust, part of the reason might also be the (often far from trivial) decoupling of the bulge from the disc.  These scaling relations exist for galaxy parameters at various radii. While photometric quantities inside one effective radius are easy to measure, the inherent limitations of traditional (single-aperture or long-slit) spectrographs have restricted the measurement of the stellar velocity dispersion to the central regions of galaxies. For instance, papers based on the Sloan Digital Sky Survey (SDSS) data \\citep[e.g.][]{bernardi03a,graves10} use velocity dispersions that have been determined from central galaxy apertures, corrected to effective velocity dispersions using standard aperture corrections. In this paper we follow the approach of \\citet[][hereafter Paper IV]{cappellari06} and make use of the panoramic capabilities of \\sauron\\ to measure velocity dispersions in circular apertures going out to an effective radius (\\se), and present scaling relations for which all the parameters are measured within the same aperture. This method offers the interesting possibility of presenting the spiral Sa galaxies in the same relations as early-type E/S0 galaxies. When measuring \\se\\ from the integrated galaxy spectrum galaxy broadening can be caused by intrinsic velocity dispersions, or by galaxy rotation. Despite this uncertainty, \\se\\ will still be a measure of the mass in a galaxy inside \\re. In addition, these velocity dispersions will not be affected by the presence of central discs, which often show low velocity dispersions \\citep[e.g.][]{fb03}. Since \\re\\ for most of our Sa galaxies is much larger than the radius inside which the galaxy bulge dominates, the scaling relations will give us information about both the bulges and the inner discs of the Sa galaxies. This paper tends to investigate both issues by combining photometry with integral-field spectroscopy for a representative sample of E to Sa galaxies, treated in a consistent manner with a homogeneous database and methods. The importance of this last point should not be overlooked, as supposedly standard parameters can vary greatly when measured by different groups. An example of this is provided by the measurement of nuclear cusp slopes \\cite[e.g.][]{ferrarese94,byun96,gebhardt96,carollo97,rest01}.  With these goals in mind, we have carried out an optical spectroscopic survey of $72$ representative nearby E/S0 galaxies and Sa galaxies to one \\re, using the custom-designed panoramic integral-field spectrograph \\sauron\\ mounted on the William Herschel Telescope, La Palma \\citep[][hereafter Paper I]{bacon01}. The \\sauron\\ representative sample was chosen to populate uniformly $M_B$--$\\,\\epsilon$ planes, equally divided between cluster and field objects \\citep[][hereafter Paper II]{dezeeuw02}. The work in this paper builds on previous results of our survey on scaling relations in other wavelength domains (Paper IV; \\citealt{jeong09}, hereafter Paper XIII). The reader is referred to other papers of the \\sauron\\ survey for results on the stellar kinematics \\citep{emsellem04} and kinematic classification \\citep{emsellem07,cappellari07} of early-type galaxies and their stellar populations \\citep{kuntschner06,shapiro10,kuntschner10} and on the kinematics \\citep{fb06} and population of early spirals \\citep{peletier07}. Hereafter, we will refer to them as Paper III, IX, X, VI, XV, XVII, VII and XI respectively.\\looseness-2  We present in this paper homogeneous ground-based $V$-band and \\textit{Spitzer} 3.6$\\mu$m-band imaging observations of the $24$ elliptical E, $24$ lenticular S0 and 24 spiral Sa galaxies of the \\sauron\\ representative sample. Aperture photometry (growth curve analysis) is carried out to homogeneously derive a number of characteristic quantities to which the more complex \\sauron\\ integral-field observations are compared. We introduce the sample selection, biases and completeness in \\S~\\ref{sec:sample}. The observations and basic data reduction are presented in \\S~\\ref{sec:obs_data}. We describe the aperture photometry and determination of the spectroscopic quantities in \\S~\\ref{sec:aper_phot} and \\S~\\ref{sec:sauron}. Bivariate scaling relations are shown in \\S~\\ref{sec:scaling}, while the Fundamental Plane relation is specifically addressed in \\S~\\ref{sec:FP}. We summarise our results and conclude briefly in \\S~\\ref{sec:conclusions}. Description of the stellar population synthesis models and methods used to derive stellar mass-to-light ratios for our galaxies are presented in Appendices~\\ref{app:models} and \\ref{app:ML}. Scaling relations showing the dependencies with kinematic substructure and environment are shown in Appendix~\\ref{app:relations}. Tables with the measured quantities are presented in Appendix~\\ref{app:photab}.\\looseness-2  Throughout the paper we adopt the \\textit{WMAP} (Wilkinson Microwave Anisotropy Probe) cosmological parameters for the Hubble constant, the matter density and the cosmological constant, of respectively $H_0=73$\\,\\kmsM, $\\Omega_M=0.24$ and $\\Omega_L=0.76$ \\citep{spergel07}, although these parameters only have a small effect on the physical scales of the galaxies due to their proximity.   \\begin{figure*} \\begin{center} \\includegraphics[angle=90,width=0.99\\linewidth]{sauron_photpars_MDM_Spitzer.eps} \\caption{Photometric characterisation of the \\sauron\\ survey. Top row: $V$-band effective radius (\\re), mean effective surface brightness (\\mue) and absolute total magnitudes. Bottom row: same quantities in the 3.6$\\mu$m band from the \\textit{Spitzer}/IRAC dataset. The sample is defined such that galaxies uniformly populate the desired ($B$-band) absolute magnitude range and that our integral-field observations reach typically as far as one \\re\\ (see \\citealt{dezeeuw02}).} \\label{fig:phot} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  In this paper we report results from photometric follow-up conducted in the context of the \\sauron\\ project. We use ground-based MDM $V$-band and \\textit{Spitzer}/IRAC 3.6$\\mu$m-band imaging to characterise our sample of E, S0 and Sa galaxies. We perform aperture photometry to derive homogeneous half-light radii, mean effective surface brightnesses and total magnitudes. Combined with the \\sauron\\ integral-field spectroscopic observations, this allows us to explore and understand the location of the galaxies in the main scaling relations as a function of the level of rotation, kinematic substructure, stellar populations and environment.  A number of conclusions can be derived from the work presented and are summarised in the following points:  \\begin{itemize}  \\item The level of kinematic substructure (as determined from our stellar kinematic maps) or the environment do not lead to a preferred location in any of the scaling relations investigated. However, this is subject to the potential biases introduced by our sample selection. Our \\textit{cluster} environment is defined by the Virgo Cluster and Leo group.  \\item The Slow Rotator (SR) galaxies define tighter relations than Fast Rotator (FR) galaxies. While this is not totally unexpected (i.e. SR are uniformly old systems), it is remarkable how the trend is kept both for small and large \\re. We find that the SR and FR  galaxies do not populate distinct locations in the scaling relations. The Sa family is the main contributor to the scatters.  \\item Sa galaxies deviate from the colour-magnitude and colour-\\se\\ relations due to the presence of dust. SR/FR galaxies instead define very tight relations.  \\item Surprisingly, extremely young objects do not display the bluest $(V-[3.6])$ colours of our sample, as is usually the case for optical colours. This can be understood in the context of the large contribution of TP-AGB stars to the infrared even for young populations, which makes the $(V-[3.6])$ colour almost insensitive to age for populations above $\\approx1$~Gyr. This effect results in a very tight $(V-[3.6])-$\\se\\ relation that allows us to define a strong correlation between metallicity and \\se.  \\item  A large number of Sa galaxies appear to follow the main relations defined by earlier-type systems, and are not necessarily offset as previously found.  We believe this is due to the use of $\\sigma_e$ (the stellar velocity dispersion integrated within \\re), which is not as much influenced by the presence of young inner discs and accounts to a large extent for rotation in the galaxies.  \\item FR and Sa galaxies that appear offset from the relations correspond mostly to objects with extremely young populations and signs of on-going, extended star formation (as already highlighted in Papers XIII and XV). For the specific case of the Fundamental Plane, we made an attempt to correct for this effect so that all galaxies are part of a tight, single relation. Once this is done, the FP coefficients in the $V$- and 3.6$\\mu$m-bands are the same within the uncertainties. The new estimated coefficients suggest that differences in stellar populations account for about 50\\% of the observed tilt with respect to the virial prediction.  \\item The observed scatter of the SR family around the Fundamental Plane is smaller than that of the FR or Sa galaxies. After correcting for stellar populations, the SR family shows almost no intrinsic scatter around the best-fit Fundamental Plane.  \\item The use of a velocity dispersion within a small aperture (e.g. \\re/8) in the Fundamental Plane results in an increase of around 15\\% in the intrinsic scatter and an average 10\\% decrease of the tilt away from the virial relation (as a result of the correlation between \\re\\ and \\se, given by the size-luminosity and Faber-Jackson relations).\\looseness-1  \\end{itemize}  We have deliberately focused on the role of stellar populations on the main scaling relations of early-type galaxies. The effects other galaxy properties exert on them (e.g. non-homology, anisotropy and dark matter variations) will be studied in ongoing and upcoming integral-field spectroscopic surveys."
    },
    "1107/1107.5311_arXiv.txt": {
        "abstract": "Using galaxy group/cluster catalogs created from the Sloan Digital Sky Survey Data Release 7, we examine in detail the specific star formation rate (SSFR) distribution of satellite galaxies and its dependence on stellar mass, host halo mass, and halo-centric radius. All galaxies, regardless of central-satellite designation, exhibit a similar bimodal SSFR distribution, with a strong break at $\\ssfr \\approx 10^{-11}\\yrinv$ and the same high SSFR peak; in no regime is there ever an excess of galaxies in the `green valley'. Satellite galaxies are simply more likely to lie on the quenched (`red sequence') side of the SSFR distribution. Furthermore, the satellite quenched fraction excess above the field galaxy value is nearly independent of galaxy stellar mass. An enhanced quenched fraction for satellites persists in groups with halo masses down to $3 \\times 10^{11}\\msol$ and increases strongly with halo mass and toward halo center. We find no detectable quenching enhancement for galaxies beyond $\\sim2\\,\\rvir$ around massive clusters once these galaxies have been decomposed into centrals and satellites. These trends imply that (1) galaxies experience no significant environmental effects until they cross within $\\sim\\rvir$ of a more massive host halo, (2) after this, star formation in active satellites continues to evolve in the same manner as active central galaxies for several Gyrs, and (3) once begun, satellite star formation quenching occurs rapidly. These results place strong constraints on satellite-specific quenching mechanisms, as we will discuss further in companion papers. ",
        "introduction": "Galaxy properties depend on their local environment. Galaxies in denser regions, such as groups and clusters, exhibit attenuated star formation rates (SFR), and therefore significantly higher red sequence fractions, as well as more elliptical morphologies as compared with galaxies in less dense environments \\citep{Oem74, DavGel76, Dre80, PosGel84}. With the advent of voluminous galaxy surveys such as the Sloan Digital Sky Survey \\citep [SDSS;][]{SDSS}, numerous works have quantified the correlations between these galaxies properties and their environment in considerable detail \\citep[e.g.,][]{HogBlaBri04, BalBalNic04, KauWhiHec04, BlaEisHog05}, and more recent works have shown that similar environmental dependence persists at least out to $z \\approx 1$ \\citep{CucIvoMar06, CooNewCoi07, PenLilKov10}.  There are two particularly interesting features of the dependence of these properties on environment. First, environmental impact on SFR (or color) is stronger than on morphology \\citep{KauWhiHec04, BlaEisHog05, BalLovBru08, BamNicBal09}: at fixed morphology, SFR still exhibits strong changes with environment, but at fixed SFR, there is almost no dependence of morphology on environment. Second, SFR/color and morphology primarily depend on small-scale ($\\lesssim 1\\mpc$) environment, with little-to-no additional dependence on environment measured on larger scales \\citep{HogBlaBri04, KauWhiHec04, BlaEisHog05}. Several works have extended this result through the use of galaxy group catalogs, showing that galaxy color and star formation history most directly depend on the properties of their host dark matter halo \\citep{BlaBer07, WilZibBud10, TinWetCon11}. Thus, understanding galaxy star formation requires identifying and understanding a galaxy's host halo properties.  The idea that environmental dependence primarily manifests itself via host halo properties is motivated by the fact that the region within a halo is physically distinct from the field environment. The halo virial radius corresponds to a physical transition between infall and virialized motions, capable of supporting strong shock fronts, at least in halos more massive than $\\sim10^{12}\\msol$ \\citep[e.g.,][]{DekBir06}. Within the virial radius, high densities lead to strong tidal forces, and accreted gas is thermalized to high temperature and pressure. `Satellite' galaxies that fall into more massive halos thus experience tidal and ram-pressure stripping, heating, and are unable to accrete gas as efficiently as in the field. Thus, within groups and clusters, it is important to distinguish satellite galaxies from the galaxy at the minimum of the host halo potential well (the `central' galaxy), which occupies a special dynamical location and is typically the oldest and most massive in the halo.  In this paper, we focus on the SFRs of satellites galaxies, how they depend on their host halo properties, and how they compare with central galaxies, which have not (yet) fallen into a more massive halo. Understanding satellite star formation is important not only for elucidating the physical processes that occur within groups and clusters, but also for a comprehensive understanding of galaxy evolution. For example, satellites are responsible for more than a 1/3 of the build-up of the red sequence population \\citep{vdBAquYan08, TinWet10}. In addition, many methods of optically identifying galaxy groups/clusters rely on selecting red sequence galaxies \\citep[e.g.,][]{KoeMcKAnn07a} because they are more prominent within groups/clusters, making it important to characterize in detail how the satellite red sequence fraction depends on halo mass and radius.  Even though it has long been known that galaxies in groups/clusters are significantly more likely to have attenuated SFRs \\citep{DreGun83, BalMorYee97, PogSmaDre99}, the mechanism(s) that dominates this transformation remains unclear, though several have been proposed. As a satellite galaxy orbits through its host halo's hot, virialized gas, ram-pressure can strip cold gas directly from the disc \\citep{GunGot72, AbaMooBow99}, causing a rapid ($\\lesssim500\\myr$) quenching of star formation. There are striking, direct examples of ram-pressure stripping occurring for satellites in nearby clusters \\citep[e.g.,][]{ChuvGoKen09}, some of which are even near the virial radius. While this suggests that ram-pressure stripping might be responsible for satellite quenching in massive clusters, even at large radii, it remains unclear whether ram-pressure can act efficiently in lower mass groups, particularly if virial shock fronts are not easily supported.  Other mechanisms affect gas accretion, both into the subhalo that surrounds a satellite and cooling onto the disc. Referred to as `strangulation' or `starvation', this more gradual process involves the lack of accretion of new hot gas, as well as loss of existing extended hot gas via heating/evaporation, tidal stripping, or ram-pressure stripping \\citep{LarTinCal80, BalNavMor00, KawMul07, McCFreFon08}, causing SFR in the disc to decline on a cold gas consumption timescale. X-ray observations show direct evidence for the truncation of hot gas halos around satellite galaxies in groups and clusters \\citep{SunJonFor07, JelBinMul08}.  Finally, galaxies in groups and clusters can interact with each other gravitationally. `Harassment' involves high-speed, nearby encounters that tidally heat satellites \\citep{FarSha81, MooLakKat98}. More violently, satellites can merge with one another \\citep{MakHut97}, a process that may not be uncommon given that satellite orbits are strongly correlated \\citep{AngLacBau09, WetCohWhi09a, WetCohWhi09b, WhiCohSmi10, Coh12}. These dynamical processes potentially could induce rapid gas consumption and thus quenching, in addition to driving morphological transformation.  To understand which of the above physical process(es) dominates, it is crucial to have a detailed understanding of how satellite SFR depends on its host halo's mass and radius within the halo, including whether there is a minimum halo mass for affecting satellite SFR. The above physical processes have different predictions for these dependences. However, the extent to which a satellite's SFR depends on its host halo mass, and how this compares with intrinsic dependence on its stellar mass, remains in debate. Several recent works using SDSS group/cluster catalogs found evidence that satellites are more likely to be quenched/red in more massive halos: \\citet{WeivdBYan06a} examining both SFR and color using the group catalog finder of \\citet{YanMovdB05a}; \\citet{BlaBer07} examining color using their own group catalog; \\citet{KimSomYi09} examining UV-based SFR using the group catalog of \\citet{YanMovdB07}; and \\citet{HanSheWec09} examining color in the MaxBCG catalog as a function of cluster richness. However, other works conclude that the satellite SFR/color dependence on halo mass is non-existent or at least much smaller than the intrinsic dependence on stellar mass: \\citet{FinBalZar08} examining SFR in the C4 cluster catalog as a function of cluster velocity dispersion; and \\citet{vdBPasYan08}b and \\citet{PasvdBMo09} examining color and SFR, respectively, using the group catalog of \\citet{YanMovdB07}. Similarly, many recent works found that satellites are more likely to be red \\citep{DePColPea04, BlaBer07, HanSheWec09} and have lower SFRs \\citep{BalNavMor00, EllLinYee01, WeivdBYan06a, vdLWilKau10, PreBalJam11} toward group/cluster center, including weak trends for groups at $z \\approx 1$ \\citep{GerNewFab07}. Though again, there is debate about the steepness of radial gradients, the extent to which they depend on halo mass, and whether they are a byproduct of satellite mass segregation \\citep{vdBPasYan08, PasvdBMo09}.  There are several possible reasons for these disagreements. First, one must compare galaxies at fixed stellar mass because SFR and color depend strongly on mass \\citep{KauHecWhi03}, and more massive galaxies are more common in denser environments \\citep{HogBlaEis03}. The environmental trends noted in many previous works were largely a reflection of these intrinsic stellar mass dependencies. Second, some SFR metrics are more susceptible to systematic biases than others, and in particular, simple color cuts can overestimate the quenched fraction because of dust reddening \\citep[e.g.,][]{MalBerBla09}. Third, in order to elucidate these trends, it is important to study satellite properties over a wide range of both satellite and host halo masses.  Finally, to fully understand the physical mechanisms operating on satellites, it is crucial to go beyond simple SFR/color cuts and understand the nature of the full SFR distribution. \\citet{BalBalNic04} examined the $u-r$ color distribution of galaxies in SDSS as a function of projected galaxy environment (defined by fifth nearest neighbor) and found that the color bimodality extended to all environments, and that the colors of blue galaxies do not depend significantly on environment. Using several different SDSS group catalogs, \\citet{Ski09} showed that the $g-r$ color distribution of satellites persists across group varying richness. More recently, \\citet{McGBalWil11} found similar trends examining the SFR distributions of all galaxies in groups and in the field both in SDSS and at $z \\approx 0.4$.  In this work, we explore the satellite galaxy SFR distribution in detail, including its dependence on host halo mass and halo-centric radius, to robustly characterize the properties of satellite star formation evolution and quenching in groups and clusters. To observationally determine the host halo masses of SDSS galaxies, we construct galaxy group catalogs using a variant of the \\cite{YanMovdB05a} group-finding algorithm, which also allows us to examine how the SFRs of satellites differ from central galaxies of similar mass. Our galaxy sample and group finder allows us to explore a considerable range of halo and satellite masses to determine the regimes in which halo-specific processes are important. We examine galaxies selected in narrow bins of stellar mass to isolate host halo dependence from intrinsic stellar mass dependence, and we use spectroscopically-derived SFR measurements which are free from dust reddening effects.  This paper represents the second in a series of four. In \\citet{TinWetCon11}, hereafter referred to as \\citetalias{TinWetCon11}, we presented our method for identifying galaxy groups in SDSS. Using the $4000$\\AA\\ break, $\\dnfk$, as a measure of recent star formation history, we showed that central and satellite galaxy star formation is nearly independent of large-scale environmental density beyond its dark matter host halo. In this work, we explore the full SFR distribution of satellite galaxies\\footnote{ We have examined all trends in this paper selecting galaxies on both $M_r$ and $\\mstel$ as well as examining both $\\dnfk$ and SSFR, and while these lead to slight quantitative differences, they do not change any of our results qualitatively.} and its dependence on stellar mass, halo mass, and halo-centric radius to fully characterize the local density dependence of satellite star formation. In \\citet{WetTinCon12a}, hereafter referred to as \\citetalias{WetTinCon12a}, we will combine the results of this paper with a high-resolution cosmological simulation that tracks satellite infall times in order to examine satellite star formation histories and quenching timescales in detail. Finally, in \\citet{WetTinCon12b}, hereafter referred to as \\citetalias{WetTinCon12b}, with the results of this paper we will use satellite orbital histories from simulation to test and constrain the physical mechanisms of satellite-specific quenching.  For all calculations we use a flat, $\\Lambda$CDM cosmology of $\\Omega_{\\rm m}=0.27$, $\\Omega_{\\rm b}=0.045$, $h=0.7$, $n_{\\rm s}=0.95$ and $\\sigma_8=0.82$, consistent with a wide array of observations \\citep[see, e.g.,][and references therein]{KomSmiDun11}. ",
        "conclusions": "\\label{sec:summary}  Using galaxy group catalogs created from SDSS Data Release 7, we examined in detail the SSFRs of satellite galaxies and how they depend on galaxy mass, host halo mass, and halo-centric radius. Our galaxy sample and group catalog provide good statistics across of wide range of masses both for satellites ($\\mstel = 5 \\times 10^{9} - 2 \\times 10^{11}\\msol$) and their host halos ($\\mthm = 3 \\times 10^{11} - 10^{15}\\msol$), and our use of spectroscopically determined SSFR means our results are insensitive to dust reddening, important for low-mass galaxies Our main results are as follows.  \\textbf{Persistent Bimodality:} All galaxies that we probe, regardless of galaxy mass, host halo mass, and halo-centric radius, exhibit similar bimodal SSFR distributions, including a fixed break at $\\ssfr = 10^{-11}\\yrinv$, the same fraction of galaxies at the break, and the same SSFR peak for active galaxies. Satellites are simply more likely to lie on the quenched side of the distribution. This persistent bimodality implies that a significant fraction of satellites have evolved in the same way as central galaxies, independent of their host halo, for several Gyrs. Once begun, the satellite SFR fading process is rapid as compared with time over which satellites remain active after infall and thus is likely shorter than a $\\sim2.4\\gyr$ cold gas consumption timescale, but it does not depend on host halo mass or halo-centric radius.  \\textbf{Clear host halo mass and radius dependences within the virial radius:} At fixed stellar mass, satellites are more likely to be quenched than central galaxies, even down to the lowest halo masses we probe ($3 \\times 10^{11}\\msol$). This quenched likelihood increases at smaller halo-centric radius, further implying that satellite quenching takes several Gyrs to occur after infall. The quenched likelihood also monotonically increases with halo mass across our entire halo mass range, even at fixed $R / \\rvir$, which implies that satellite quenching does not simply depend on \\textit{current} local density. It is not immediately clear whether satellite quenching is more efficient/rapid in higher mass halos, or surviving satellites in that regime have simply been satellites longer.  \\textbf{Unclear stellar mass dependence:} The central galaxy quenched fraction increases significantly with stellar mass. While more massive satellites show a weaker absolute change in quenched fraction with host halo mass or radius, this is largely a manifestation of central galaxy mass dependence prior to infall. The satellite quenched fraction excess, which more directly measures quenching efficiency, shows little-to-no dependence on galaxy mass.  \\textbf{Importance of satellite-specific evolution:} In \\citetalias{TinWetCon11}, we showed that the star formation histories and quenched fractions of central galaxies at fixed galaxy (or halo) mass are essentially independent of their large-scale ($10\\hmpc$) environment. Here, we have shown that this lack of environmental dependence extends to central galaxies as close as $\\sim2\\,\\rvir$ to massive clusters. Together, these results demonstrate the importance of satellites in understanding galaxy evolution, because satellite-specific processes are the only significant environmental processes that affect galaxy star formation and color.  The results of the satellite quenched fraction excess suggest that infalling active satellites quench at least as efficiently/rapidly at high stellar mass as at low stellar mass. However, it remains true that satellite-specific processes are less important for building up the red sequence at higher stellar mass, because more galaxies quench as central galaxies prior to infall in that regime \\citep[see also][]{vdBAquYan08}. Unfortunately, it is difficult to determine directly from $\\fsatqexcess$ whether satellite quenching is fully independent of galaxy mass because doing so requires knowing in detail the quenched fraction as a function of stellar mass prior to satellite infall. In \\citetalias{WetTinCon12a}, we will use a high-resolution cosmological simulation to obtain the infall time distributions of satellites, combined with observational constraints on high-redshift quenched fractions, to examine this issue more carefully.  The dependences on mass and halo-centric radius that we see broadly agree with many previous works using SDSS catalogs. This includes a monotonic increase in quenched fraction with halo mass mass which nearly equals the intrinsic dependence on galaxy mass \\citep{WeivdBYan06a, BlaBer07, HanSheWec09, KimSomYi09}, as well as satellites being more likely to be quenched at smaller radii \\citep{WeivdBYan06a, BlaBer07, HanSheWec09} with less massive satellites exhibiting stronger radial gradients \\citep{vdLWilKau10}. Our catalog enabled us to extend many of these results to a wider range of galaxy and halo masses, particularly at the low mass end, and our use of spectroscopically-derived SSFR ensures that of quenched fraction results are not biased by dust in that regime. Furthermore, our SSFR bimodality results qualitatively agree with previous results on the color bimodality as a function of projected galaxy density \\citep{BalBalNic04, BalBalBow06} and the more recent results of \\citet{McGBalWil11} and \\citet{PenLilRen11}, who showed similar SSFR distributions and active galaxy SSFR values when splitting all galaxies into groups vs. the field. We have shown that this `universal' SSFR bimodality extends to specifically satellite galaxies in all halo masses and at all radii.  Using color cuts to measure red galaxy fractions, \\citet{vdBAquYan08}, and more recently \\citet{PenLilRen11}, employed a similar statistic as our satellite quenched fraction excess to show that this excess is independent of satellite mass. We have shown that this results extends to more robust measures of SFR, and to all halo masses and at all radii. However, our physical interpretation of this statistic differs from theirs. In particular, this statistic only measures the \\textit{excess} fraction of satellites that quenched after infall that would not have quenched had they remained central galaxies, because the $z \\approx 0$ central galaxy quenched fraction does not represent the initial conditions of satellites at the time of infall. We will explore this issue further in \\citetalias{WetTinCon12a}.  The mass and radius dependences we see disagree with those of \\citet{vdBPasYan08}b, who found only weak dependence of galaxy properties on halo mass and radius at fixed stellar mass. However, they examined primarily mean $g-r$ color of satellites as a measure of star formation, which is problematic because dust reddening causes a severe biasing of $g-r$ color, particularly for low-mass galaxies, and the mean value is a less sensitive characterization of a highly bimodal distribution. Indeed, we find similar results to theirs when examining mean $g-r$ color.  The satellite SSFR distribution dependencies we have detailed in this paper provide tight constrains on models for satellite-specific evolution. In \\citetalias{WetTinCon12a}, we will use a high-resolution cosmological simulation to obtain the detailed orbital histories of satellites, which we will use to more rigorously constrain the timescales of satellite quenching, and in \\citetalias{WetTinCon12b} we will develop models for satellite-specific quenching via different mechanisms to test which reproduce the trends observed in this paper."
    },
    "1107/1107.5902_arXiv.txt": {
        "abstract": "The central star-forming region in a blue compact dwarf galaxy, II Zw 40, was observed in the 340~GHz ($880~\\micron$) band at $\\sim 5$ arcsec (250 pc) resolution with the Submillimetre Array (SMA). A source associated with the central star-forming complex was detected with a flux of $13.6\\pm 2.0$ mJy. The structure is more extended than the beam in the east--west direction. The SMA 880 $\\micron$ flux is analyzed by using theoretical models of radio spectral energy distribution along with centimetre interferometric measurements in the literature. {We find (i) that the SMA 880 $\\micron$ flux is dominated ($\\sim 75$ per cent) by free--free emission from the central compact star-forming region, and (ii) that the contribution from dust emission to the SMA 880 $\\micron$ flux is at most $4\\pm 2.5$ mJy. We also utilize our models to derive the radio--FIR relation of the II Zw 40 centre, suggesting that free--free absorption at low frequencies ($\\nu\\la$ several GHz; $\\lambda\\ga$ several cm) and spatial extent of dust affect the radio--FIR relation.} ",
        "introduction": "Blue compact dwarf galaxies (BCDs) generally host compact and ongoing star formation activities in metal-poor and gas-rich environments \\citep*{sargent70,vanzee98}. Because both low metallicity and rich gas content indicate an early evolutionary stage, BCDs can be used as nearby laboratories of primeval galaxies which should exist at high redshift. In some BCDs, the most active class of star formation is taking place in super star clusters (SSCs) {\\citep{turner98,kobulnicky99,johnson03}, which are expected to `mimic' the starburst in high-redshift primeval galaxies}.  II Zw 40 is a well studied BCD with a low oxygen abundance, $12+\\log (\\mathrm{O/H})=8.13$ \\citep{thuan05}. This galaxy hosts a high star formation activity associated with the SSCs in the centre. At centimetre wavelengths, the central star-forming region in II~Zw~40 is compact and optically thick for free--free absorption, and is categorized as a `supernebula' or `ultradense \\hii\\ region' \\citep{turner98,kobulnicky99}. Such a dense and compact star formation activity is called `active' mode in \\citet{hunt03} and \\citet{hirashita04}. II Zw 40 is also classified as a Wolf-Rayet galaxy: the Wolf-Rayet feature indicates that the typical age of the current starburst is a few Myr \\citep{vacca92}. \\citet{buckalew05} derived an age of 2.6 Myr from the H$\\alpha$ and H$\\beta$ equivalent widths. The Br$\\gamma$ equivalent width also shows age $\\la 3$ Myr \\citep{vanzi08}. Stellar spectral synthetic models support young ages $\\sim 2$ Myr, although there is an underlying old stellar population \\citep{westera04}.  In order to trace dense `embedded' star-formation activities, optically thin star formation indicators are useful, such as far-infrared (FIR) dust luminosity \\citep[e.g.][]{kennicutt98,inoue00} and radio (thermal plus non-thermal) luminosity \\citep{condon92}. Indeed, there is a correlation between FIR and radio luminosities in nearby star-forming galaxies \\citep[e.g.,][]{dejong85,helou85}, which is naturally explained if both luminosities are strongly connected with star formation activities \\citep{volk89}. However, the studies of radio--FIR relation are biased to objects with FIR detection \\citep[e.g.][]{condon92}, which means that a significant dust enrichment has already occurred. Therefore, the evolution of radio--FIR relation in young primeval galaxies is not yet clear. The gas density, the magnetic field strength, and the energy density of cosmic ray electrons affect the evolution of radio emission on various time-scales {\\citep{helou93,niklas97,murphy09,lacki10}}, while the dust enrichment plays a role to increase the FIR luminosity \\citep{hirashita_hunt08}.  There have been some observational studies on the evolution of radio--FIR relation along the cosmic age. The radio--FIR relation at moderate and high redshifts ($z\\la 4$) is broadly similar to that at the local Universe \\citep*[e.g.][]{garrett02,gruppioni03,ibar08, murphy09,michalowski10}, although there is also a slight indication of evolution \\citep{vlahakis07,seymour09,michalowski10}. The disadvantage of high-$z$ observations is the difficulty in deriving quantities related to galaxy evolution (age, metallicity, etc.) with high accuracy. On the other hand, studies of nearby metal-poor dwarf galaxies as `laboratories' of primeval galaxies provide an alternative way to approach the properties of primeval galaxies. \\citet*{hopkins02} and \\citet{wu08} conclude that the radio--FIR relation of star-forming dwarf galaxies is similar to that of normal galaxies in spite of the difference in metallicity \\citep*[see also][]{klein91}, although \\citet{cannon05,cannon06} find a significant deviation from the canonical radio--FIR relation for some individual dwarf galaxies.  In this paper, we do not take a statistical way, but investigate a single object in details. II Zw 40 is suitable for studying radio--FIR emission in a metal-poor young object, since as mentioned above it hosts an extremely young ($\\la 3$ Myr) active star formation with a metallicity of 1/4 Z$_{\\sun}$ (for the solar metallicity, we adopt $12+\\log\\mathrm{O/H})=8.69$; \\citealt{lodders03}). In order to spot the young star-forming component, a high resolution is necessary. The star-forming region is resolved well at centimetre wavelengths, while there is no information on FIR--submillimetre (submm) dust emission on such a small scale. A submm interferometric observation is desired for the purpose of resolving the dust emission in the star-forming region. Therefore, we observed II Zw 40 by the Submillimetre Array (SMA; \\citealt{ho04}), and we report on this observation in this paper. This observation, combined with centimetre radio interferometric data in the literature, will enable us to obtain the radio--FIR emission properties (or radio--FIR relation) in the very young star-forming region with a low metallicity.  This paper is organized as follows. We explain the observations and the data reduction in Section \\ref{sec:obs}, and describe observational results in Section \\ref{sec:result}. In Section \\ref{sec:model}, we interpret the results along with the radio data in the literature by using theoretical models. In Section \\ref{sec:firsed}, we analyze the contribution from dust emission to the SMA flux. We discuss the radio--FIR relation based on the SMA observation in Section \\ref{sec:discuss}. Finally, we conclude in Section \\ref{sec:conclusion}. The distance to II Zw 40 is assumed to be $D=10.5$ Mpc ($cz=789$ km s$^{-1}$ with $H_0=75$ km s$^{-1}$ Mpc$^{-1}$). At this distance, 1 arcsec corresponds to 50.9 pc. ",
        "conclusions": "\\label{sec:conclusion}  In order to reveal the radiative properties of young active starburst, the central star-forming region in II Zw 40 was observed in the 340~GHz (880 $\\micron$) band at $\\sim 5''$ resolution with SMA. A source associated with the central star-forming complex was detected with a flux of $13.6\\pm 2.0$ mJy and a 10-arcsec elongation in the east--west direction. The flux in the central part of II Zw 40 has been analyzed by using the theoretical radio SED model developed by \\citet{hirashita06}, and interpreted along with interferometric measurements at centimetre wavelengths in the literature. Then, we have found { \\begin{enumerate} \\item that the SMA 880 $\\micron$ flux is dominated by free-free emission, and \\item that possible contribution from dust emission to the SMA flux is less than $4\\pm 2.5$ mJy. \\end{enumerate} } Our models have been used to derive the radio--FIR relation of the II Zw 40 centre. We have discussed free--free absorption at low frequencies ($\\nu\\la$ several GHz; $\\lambda\\ga$ several cm) and spatial distribution of dust ($R_\\mathrm{dust}$) as possible factors affecting the radio--FIR relation."
    },
    "1107/1107.5243_arXiv.txt": {
        "abstract": "Covariant density functional theory is used to study the effect of strong magnetic fields, up to the limit predicted for neutron stars (for magnetars $B \\approx10^{18}$G), on nuclear structure. All new terms in the equation of motion resulting from time reversal symmetry breaking by the magnetic field and the induced currents, as well as axial deformation, are taken into account in a self-consistent fashion. For nuclei in the iron region of the nuclear chart it is found that fields in the order of magnitude of $10^{17}$G significantly affect bulk properties like masses and radii. ",
        "introduction": "Several studies (e.g. \\cite{Shapiro1983,Haensel2006}) have determined the presence of intense magnetic fields, up to $10^{16}$ G, on the surface of neutron stars.  Theoretical models suggest that these magnetic fields might reach up to $B \\approx10^{18}$ G, and even larger values if one considers the limit imposed by the virial theorem ($B \\approx2 \\cdot10^{18}$ G \\cite{Broderick2000}). The influence of the magnetic fields on the equation of state (EOS) has been thoroughly studied and reported in last years (e.g. \\cite{Broderick2000,Guang-Jun2003,Wei2006,Rabhi2008}). It should be noted lower than $10^{18}$~G magnetic fields can influence the low density parts of a neutron star, such as its surface layer.   The outer crust is fundamentally composed of well separated nuclides, and its structure determined by the energies of isolated nuclei, the kinetic energy of electrons and the lattice energy. Thus its composition depends very much on the binding energy of stable and unstable nuclei in the outer crust below the neutron drip density and of neutron-rich nuclear systems above the neutron drip. At the lowest densities it is thought that the most abundant component is $^{56}$Fe because of its high binding energy, a fact actually observed in emission spectra (e.g. \\cite{Cackett2008}).  The magnetic field strength required to directly influence the EOS can be estimated by considering its effects on charged particles. Charge-neutral, beta-equilibrated neutron star matter contains both negatively charged leptons (electrons and muons) and protons. Magnetic fields quantize the orbital motion (Landau quantization) of these charged particles. When the Fermi energy of the proton becomes significantly affected by the magnetic field, the composition of matter in beta equilibrium is modified. This is reflected in a change of the pressure of matter. It has been found in Ref.~\\cite{Broderick2000} that this occurs for fields of approximately $10^{18}$ G, and that in general leads to a stiffening of the EOS.  However, there are very few studies \\cite{Kondratyev2000,Kondratyev2001} of the changes that these very intense magnetic fields may eventually bring to the composition of the crust. The structure and composition of the crust is important in the thermal and rotational evolution of neutron stars, in particular in the theory behind glitches \\cite{Shapiro1983}. Some other studies \\cite{Duncan1997,Thompson2001} point to a magnetically driven crust activity as the source of soft gamma repeaters (SGR).  So far, the impact of intense magnetic fields in nuclei found in the outer crust has been studied on a qualitative way using a simple non self-consistent method by Kondratyev and collaborators \\cite{Kondratyev2000,Kondratyev2001}. It was found that fields as low as $10^{16}$ G may modify the nuclear shell structure, well within the range of theoretically possible magnetic strengths. There are, however, still some questions that need to be addressed: i) What is the minimum field that is able to significantly alter the nuclear structure? ii) Is this field low enough to be found in a significant proportion of neutron stars or magnetars?  iii) Is this effect big enough to influence astrophysically relevant situations and processes, e.g. neutron star outer crust composition or final element abundances in nucleosynthesis scenarios?  The objective of the present work is to try to find answers to these questions, in a quantitative way if possible, using a fully microscopical description of the nuclear system within covariant density functional theory. The formalism used will be introduced in Section II. A general discussion of the effects of the external magnetic field on nuclei will be given on Section III, and particularized to an example nucleus. A discussion of the possible influence of the magnetic fields on neutron star outer crust nuclei can be found on Section IV. Finally, Section V is devoted to the conclusions. ",
        "conclusions": "The influence of strong magnetic fields on nuclear structure has been studied using a fully self-consistent covariant density functional. It has been found that a field strength of at least $10^{17}$ G is needed to appreciably modify the nuclear ground state. For sufficiently high fields these effects cannot be studied using a frozen-field approach since there are level rearrangements, causing spin-polarization and induced currents and thus a self-consistent model is required. It is the advantage of covariant theories that the these currents can be taken into account without any additional parameters. The minimum magnetic field that changes the nature of the nuclear ground-state is very much dependent on the nucleus, and its effects on the nuclear binding energy per nucleon range from a few tenths of keV for $B\\approx0.5\\cdot10^{17}$G to a few hundreds of keV for the maximum field theoretically possible $B \\approx10^{18}$ G.  No neutron star has yet been observed with such intense magnetic fields $B>10^{16}$ G, even though theoretical models hint that such objects exist \\cite{Duncan1997,Thompson2001}. In such a case, the composition of its outer crust might be radically different from that of normal neutron stars. The relevance of this change in composition depends on the abundance of normal neutron stars compared with that of magnetars. Until now, observations suggest that magnetars are not so common in the universe and thus it is unlikely that the magnetic field effects on nuclear structure play an important role in global astrophysical observables like element abundances. However, changes in the composition of magnetars might be relevant in the study of different phenomena in these particular kind of neutron stars, for example elastic properties of the crust \\cite{Pethick1998}, pulsar glitches \\cite{Mochizuki1995}, or cooling \\cite{DeBlasio1995}. With the inclusion of a proper pairing interaction, the covariant DFT model presented in this work can be used to perform a quantitative exploration of all these questions. A systematic study of changes in the composition of the outer crust will be presented in a upcoming publication.  \\vskip 0.5cm  \\noindent\\textbf"
    },
    "1107/1107.5596_arXiv.txt": {
        "abstract": "Gamma ray bursts can potentially be used as distance indicators, providing the possibility of extending the Hubble diagram to redshifts $\\sim7$. Here we follow the analysis of \\citet{s07}, with the aim of distinguishing the timescape cosmological model from the \\lcdm\\ model by means of the additional leverage provided by GRBs in the range $2\\lesssim z\\lesssim7$. We find that the timescape model fits the GRB sample slightly better than the \\lcdm\\ model, but that the systematic uncertainties are still too little understood to distinguish the models. ",
        "introduction": "The timescape (TS) model is an inhomogeneous cosmological model that explains the apparently accelerated cosmic expansion first observed in the supernova luminosity distances as an artifact of gradients in gravitational energy between gravitationally bound systems and the intervening negatively curved voids. In the dynamical spacetime of general relativity, this leads to a variance in the calibration of the clock rates of ideal observers who fit average smoothed-out geometries to the underlying inhomogeneous matter distribution. The TS model agrees closely with the \\lcdm\\ model over the range of scales probed by the supernova data~\\citep{LNW}, with certain qualifications: parameter values obtained by minimizing $\\chisq$ fits to the TS Hubble curve depend significantly on the process used to reduce the \\sn\\ light curves~\\citep{sw}.  The current state of knowledge of systematic uncertainties in the \\sn\\ data precludes discrimination between the TS and \\lcdm\\ models using \\sne\\ \\citep{sw}. In fact, calculation of the effective comoving distance $H_0D(z)$ shows that in the redshift range probed by \\sne\\ there is little to distinguish between the TS model with the best-fit value for the present void fraction $\\fvn=0.762$ from the Gold dataset of \\citet{gold}. \\citet{obs} has noted that over different redshift ranges $H_0D(z)$ for the TS model closely approximates $H_0D(z)$ for spatially flat \\lcdm\\ models with different values of \\Om\\ and $\\Omega_{\\Lambda0}$. It is thus seen to interpolate between different \\lcdm\\ models as the redshift is varied (see Fig.~\\ref{fig:hod}). Fig.~\\ref{fig:hod} shows that between $z\\simeq2$ and $z\\simeq6$, the TS $H_0D(z)$ crosses from coinciding closely with the best fit line from the \\sne\\ only to that predicted by the best fit to WMAP, BAO and the \\sne. In principle, Gamma Ray Bursts (GRBs), which probe this redshift range, could distinguish the TS and \\lcdm\\ models in this redshift range, although their use as distance indicators is far from established.  This paper will establish that the TS model is also supported by the current GRB data~\\citep{s07}, but that, as one might expect from the \\sne\\ results, the uncertainties in the data are as yet too large to distinguish the models in the redshift range $2<z<6$.  The paper is organized as follows. Section~\\ref{background} explains the method of ``standardizing'' the GRBs for their use as distance indicators, and gives a brief derivation of the TS luminosity distance. Section~\\ref{results} describes the results, before a discussion and conclusion are presented in Section~\\ref{discussion}.  \\begin{figure} \\begin{center} \\caption{Effective comoving distance as a function of redshift for various spatially flat models (dotted lines) and for the TS model with $\\fvn=0.762$ (solid line). Parameter values for the dotted lines are (i) $\\OmMn=0.249$ (best-fit to WMAP only); (ii) $\\OmMn=0.279$ (joint best-fit to WMAP, BAO and \\sne); (iii) $\\OmMn=0.34$ (best-fit to \\citet{gold} \\sne\\ only). After \\citet{obs}. \\label{fig:hod}} \\includegraphics[scale=0.4]{h0d.eps} \\end{center} \\end{figure}   \\section[]{Background}\\label{background} \\subsection[]{The timescape model} In keeping with current observations that the large-scale cosmic structure consists of voids of average diameter $\\sim30~h^{-1}$ Mpc~\\citep{hoyle2004,pan2011} separated and threaded by walls and filaments containing clusters of galaxies, the timescape model is based on a differentiation of the Universe into gravitationally bound spatially flat wall regions and negatively curved voids. There is a consequent small backreaction ($\\leq 5\\%$ as a normalized energy density) which nevertheless leads to significant cosmological effects over cosmological timescales~\\citep{clocks}. At late epochs the construction of a single smoothed-out geometry becomes problematic when the underlying geometry varies. The TS model is based on the assumption that different equivalent descriptions of a smoothed-out average geometry can be given, but these descriptions will vary between canonical observers who each assume that the average geometry has the same spatial curvature as the locally determined geometry. Differences in the calibration of rulers and clocks grow cumulatively as the variance in spatial geometry grows, and these must be taken into account when reconstructing the expansion history of the universe from information on null geodesics.  As observers in galaxies, our local average geometry, assumed to be spatially flat with scale factor $a_{\\rmn{w}}$, is given by \\be \\label{eq:wallmetric} \\d s^2_{\\finfty}=-\\d\\tau^2+a_{\\rmn{w}}^2(\\tau)[\\d\\eta_{\\rmn{w}}^2 + \\eta_{\\rmn{w}}^2\\d\\Omega^2]. \\ee Finite infinity~\\citep{fit1}, denoted by $\\finfty$, demarcates the boundary between gravitationally bound and unbound systems~\\citep{clocks}. A similar expression defines the negatively curved geometry at the centre of a void, with an appropriate void time parameter $\\tau_{\\rmn{v}}$ and scale factor $a_{\\rmn{v}}$.  The volume averaged scale factor \\be\\label{volumeaveragescalefactor}\\bar{a}^3 = f_{\\rmn{v}_i}a_{\\rmn{v}}^3 + (1-f_{\\rmn{v}_i})a_{\\rmn{w}}^3,\\ee where $f_{\\rmn{v}_i}\\ll1$ is the initial void fraction, evolves according to an averaging of the Einstein equations for an inhomogeneous dust cosmology due to \\citet{Buchert2000}. The corresponding averaged geometry, in terms of the proper time $t$ at a volume average position in freely expanding space, has the form \\be \\label{eq:baremetric} \\d s^2=-\\d t^2+\\bar{a}^2(t)\\d\\bar{\\eta}^2 + A(\\bar{\\eta},t)\\d\\Omega^2, \\ee where the area function $A$ is defined by an average over the particle horizon volume~\\citep{clocks}. The local time in a wall is related to $t$ by the phenomenological lapse parameter $\\bar{\\gamma}=\\frac{\\rmn{d}t}{\\rmn{d}\\tau}$.  A single geometry that relates metrics~(\\ref{eq:wallmetric}) and (\\ref{eq:baremetric}) is constructed by matching the radial null geodesics of wall and volume average geometries sharing a common centre~(\\citet{clocks}, $\\S$~5.2). Along the radial null geodesics, the line elements of the two geometries are simply related by a conformal factor. Once the metric is extended to cosmological scales, instead of (\\ref{eq:baremetric}) wall observers describe the large-scale universe by the effective metric \\be \\label{eq:dressedmetric} \\d s^2=-\\d\\tau^2+\\frac{\\bar{a}^2}{\\bar{\\gamma^2}}[\\d\\bar{\\eta}^2 + r_{\\rmn{w}}^2(\\bar{\\eta},t)\\d\\Omega^2], \\ee where $r_{\\rmn{w}}\\equiv\\bar{\\gamma}(1-f_{\\rmn{v}})^{1/3}(1-f_{\\rmn{v}_i})^{-1/3}\\eta_ {\\rmn{w}}(\\bar{\\eta},t)$. Metric~(\\ref{eq:dressedmetric}) is the ``dressed'' geometry that arises when we attempt to reconcile our position as observer within a finite infinity region with observations of objects at cosmological distances. In general, observers (localised within finite infinity regions) cannot assume that their measurements of cosmological parameters correspond to the global average values. However, due to the existence of a scale of statistical homogeneity, the values of locally measured cosmological parameters should converge towards their global average values as the averaging volume increases.  Volume average (``bare'') parameters, referred to metric~(\\ref{eq:baremetric}), differ from dressed parameters. The dressed matter density is $\\OmMn=\\bar{\\gamma}^3\\OMM$, and the Hubble parameter of metric~(\\ref{eq:dressedmetric}) is related to the bare Hubble parameter according to \\be \\label{eq:hubble} H=\\bar{\\gamma}\\bar{H}-\\frac{\\d}{\\d t}\\bar{\\gamma}. \\ee It is a feature of the TS model that the variance of parameters calculated from observations of nearby objects (on scales $\\lesssim 100~h^{-1}~$Mpc) will be relatively large because such objects lie within the scale of statistical homogeneity. Since the volume of space is dominated by voids --- of typical diameter $\\sim30~h^{-1}~$Mpc --- which appear to expand faster than the walls, observers in a typical galaxy looking on a typical line of sight through local voids will infer that their local universe is expanding faster than the global average. Eventually a typical line of sight will intersect a sufficient number of walls as well as voids to approach the global average, which does not change by sampling on ever larger scales. The transition scale from from large to small variance in the expansion must be larger than the diameter of the dominant voids, and is referred to as the \\emph{scale of statistical homogeneity}. It is expected to be comparable to the BAO scale, $\\sim 100~h^{-1}~$Mpc \\citep{equiv}. Thus any observer in a galaxy will typically see a ``Hubble bubble'' on scales $\\lesssim 100~h^{-1}~$Mpc.  In terms of volume average time, $t$, the luminosity distance for wall observers is~\\citep{sol} \\be \\label{eq:tslumdist1} d_L =\\bar{a}_0(1+z)r_{\\rmn{w}}. \\ee The Buchert equations have an exact general solution which admits a particular late-time attractor solution, to which the general solution converges to within $1\\%$ by redshift $z\\sim37$~\\citep{sol}. For this tracker solution, eq.~(\\ref{eq:tslumdist1}) becomes~\\citep{obs} \\be \\label{eq:tslumdist} \\bar{H}_0d_L&=&(1+z)^2(\\bar{H}_0t)^{2/3}\\int_{t}^{t_0}\\frac{2\\bar{H}_0\\rmn{d}t}{ (2+f_ { \\rmn { v } } (t '))(\\bar{H}_0t)^{2/3}}\\\\\\nonumber &=&(1+z)^2y^2\\times\\\\\\nonumber & &\\hspace{-0.5cm}\\Big[2y+\\frac{b}{6}\\ln\\Big(\\frac{(y+b)^2}{y^2-by+b^2}\\Big)+\\frac{b } { \\sqrt { 3 } } \\tan^ { -1}\\Big(\\frac{2y-b}{\\sqrt{3}b}\\Big)\\Big]^{y_0}_y \\ee where $y^3\\equiv\\bar{H}_0t$ and $b^3\\equiv2(1-f_{\\rmn{v}0})(2+f_{\\rmn{v}0})/(9f_{\\rmn{v}0})$. In Fig.~\\ref{fig:hod}, $D=d_L/(1+z)$, and the dressed Hubble constant $H_0$, is related to the bare Hubble constant $\\bar{H}_0$ by \\be\\label{eq:dressedH0} H_0=(4f_{\\rmn{v}0}^2+f_{\\rmn{v}0}+4)/[2(2+f_{\\rmn{v}0})]\\bar{H}_0. \\ee  The dressed Hubble constant is the one whose value should coincide with that which is conventionally determined on scales greater than the scale of statistical homogeneity, $z\\gtrsim0.033$. Using the ``Gold'' \\sn\\ dataset of \\citet{gold}, \\citet{LNW} found the dressed Hubble constant to be $H_0= 61.7^{+1.4}_{-1.3}\\kmsmpc$, (and the bare Hubble constant $\\bar H_0=48.2\\pm2.6\\kmsmpc$). However, since the supernova data magnitudes depend on an overall normalization determined from the local distance ladder, they cannot be used to determine the Hubble constant alone. Joint estimates of $H\\Z0$ and $\\fvn$ which fit both the angular diameter distance of the sound horizon in the CMB anisotropy data and the baryon acoustic oscillation scale in galaxy clustering statistics do provide an independent constraint on the value of $H\\Z0$, however, and on this basis \\citet{LNW} find the range of values of the dressed Hubble constant to be roughly constrained to lie in the interval $57\\lsim H\\Z0 \\lsim 68\\kmsmpc$.  This is lower than the S$H\\Z0$ES best estimate of \\citet{shoes} but that survey relies on the calibration of the distance ladder using objects that lie within the scale of statistical homogeneity, and this may involve complicated systematics given that higher values of $H\\Z0$ are expected below the statistical homogeneity scale. Estimates of $H\\Z0$ which do not rely on calibration with nearby objects are often somewhat lower. For example, \\citet{courbin10} find $H_0= 62^{+6}_{-4}\\kmsmpc$ using the time delay from strong gravitational lensing of quasars, and \\citet{beutler11} estimate $H_0=67\\pm3.2\\kmsmpc$ using the WMAP sound horizon-calibrated BAO signal in the 6dF galaxy survey. These measurements indicate that a value of $H_0$ consistent with the TS model is still to be obtained once systematic errors on distance determinations are reduced.  With the tracker solution, the bare densities $\\OMM$ and $\\OMk$ can be written in terms of the present void fraction. In particular, the dressed matter density, measured by wall observers, for which the numerical value is most likely to be similar to that of a FLRW model, will be \\be \\label{eq:fv0Om} \\Omega_{\\rmn{m}0} = \\frac{1}{8}(2+f_{\\rmn{v}0})^3\\OMM=\\frac{1}{2}(1-f_{\\rmn{v}0})(2+f_{\\rmn{v}0}) \\ee for the tracker solution.  Since there is no nearby GRB sample, there is no GRB calibration of $\\bar{H}_0$, and we work with relative distances only. The fitting process therefore results in a best-fit value for the single parameter $f_{\\rmn{v}0}$.   \\subsection{GRB data reduction method} In this paper I will use the sample of 69 GRBs selected by \\citet{s07} (henceforward S07) as having sufficient light curve data to compute their placement on a Hubble diagram.  GRBs are not standard candles, since their luminosities span several orders of magnitude (whether one assumes collimated or isotropic emission). However, there are ongoing attempts to ``standardize'' GRBs given their promise for cosmology: they occur at higher redshifts than any established standard candles, and radiation in the gamma band ($\\geq10$ keV) is not subject to the same limitations due to dust extinction as the optical band~\\citep{ggf2006}. Certain GRB light curve parameters have been found to correlate with each other, offering the possibility of computing a magnitude, much in the same way as the Phillips stretch-luminosity relation is used to reduce scatter in the \\sn\\ Hubble diagram.  \\citet{s07} uses four light curve parameters that correlate with the luminosity: (1) the lag time $\\lagt$ is the time shift between the hard and the soft light curves; (2) the light curve variability $V$ is the normalized variance of the light curve around a smoothed version of that light curve; (3) the peak energy $\\Epeak$ is the photon energy at which the $\\nu F_{\\nu}$ spectrum is brightest; and (4) the minimum rise time $\\trt$ is the shortest time over which the light curve rises by half the peak flux of the pulse. {\\bf A} fifth correlation, $\\Epeak-\\Egamma$, relates the peak energy of the light curve to total photon energy emitted by the burst. This is the tightest of the correlations~\\citep{g2004}, but it requires measurement of a jet break time by which the measured (isotropic-equivalent) energy can be corrected for the collimation. Along with these luminosity indicators a peak flux $P$ is also measured for a wide range of bandpasses. A bolometric flux (or fluence) can be calculated by extrapolating to high and low energies using the well-known broken power law of \\citet{band93} for the GRB spectrum and integrating over all energies. This brings consistency to the brightnesses, and given a cosmological model, permits the calculation of an isotropic luminosity \\be \\label{eq:dl} L=4\\pi d_L^2\\Pbolo. \\ee  The algorithm goes as follows. The luminosity indicator is the independent variable, and from eq.~(\\ref{eq:dl}) is obtained the $Y$-coordinate. A linear fit to the logarithms of these quantities gives an empirical relationship between the luminosity indicator and the luminosity. For this fit, we use the bisector of the two ordinary least squares fits: that of $X$ against $Y$, and then vice versa~\\citep{isobe90}. We can then use this relationship to calculate a theoretical luminosity curve for each indicator, based on the luminosity distance obtained from eq.~(\\ref{eq:dl}). In cases where a jet break has been measured, $\\Epeak$ is related to the collimation-corrected energy $\\Egamma$, by \\be \\label{eq:dl2} \\Egamma=4\\pi d_L^2\\Sbolo(1-\\cos\\thetajet)(1+z)^{-1}, \\ee for jet opening angle $\\thetajet$ and bolometric fluence $\\Sbolo$. The uncertainties in the $Y$-axis quantities $\\log~L$ and $\\log~\\Egamma$ are obtained from the uncertainties in the $X$-axis quantities in the standard way. Because the physics of the GRB explosions is not completely understood, the correlations contain some scatter over and above the measurement noise. To account for this, an additional intrinsic uncertainty is estimated such that the reduced $\\chisq$ of the indicator-luminosity calibration curve is unity. The best-fit lines for these relations are given along with their uncertainties in Appendix~\\ref{app:1}.  From each calculated $L$ or $\\Egamma$ we then recalculate a luminosity distance via (\\ref{eq:dl}) or (\\ref{eq:dl2}), from which we obtain a distance modulus in the standard way: $\\mu=5\\log d_L - 25$ for $d_L$ in Mpc. The propagated uncertainties are \\citep{s07} \\be \\label{eq:pbolounc} \\sigma_{\\mu}^2=(2.5\\sigma_{\\log L})^2+\\Big(\\frac{1.086\\sigma_{\\Pbolo}}{\\Pbolo}\\Big)^2, \\ee if the bolometric flux $\\Pbolo$ is used, or, if the bolometric fluence is used, \\be \\label{eq:sbolounc} \\sigma_{\\mu}^2=(2.5\\sigma_{\\log \\Egamma})^2+\\Big(\\frac{1.086\\sigma_{\\Sbolo}}{\\Sbolo}\\Big)^2+\\Big( \\frac{1.086\\sigma_{\\Fbeam}}{\\Fbeam}\\Big)^2, \\ee where the beam factor $\\Fbeam\\equiv(1-\\cos\\thetajet)$ is calculated from the jet break time $\\tjet$.  Finally, we take a weighted average of all the five different distance moduli: \\be \\label{eq:muav} \\mu=\\frac{\\Sigma_i\\mu_i/\\sigma_{\\mu_i}^2}{\\Sigma_i\\sigma_{\\mu_i}^{-2}}, \\ee with the uncertainty \\be \\label{eq:muaverr} \\sigma_{\\mu}=\\Big(\\Sigma_i\\sigma_{\\mu_i}^{-2}\\Big)^{-1/2}. \\ee  We avoid circularity by performing a simultaneous fit of both the cosmology and the luminosity relations~\\citep{g2004,s07}--- i.e. the luminosity relations are part of the model. The value of $H_0$ here is arbitrary, since its variation changes the Hubble line and the luminosity calibration of the data in an identical way, resulting merely in a change in the overall normalization of the Hubble diagram. GRBs do not occur in the local universe, so calibrating the GRB Hubble diagram to a value of $H_0$ with any accuracy is not possible. This is different to the \\sn\\ case, in which the calibration of light curve and stretch can be done model-independently with nearby \\sne, and then extrapolated to objects at higher redshifts. However, regardless of the normalization, the \\emph{shape} of the curve in the Hubble diagram depends solely on $\\fvn$. This means that for a range of values of $\\fvn$, here between 0.0 and 1.0, we calibrate the luminosity relation and compute the placement of the GRBs on the Hubble diagram, and calculate a corresponding range of $\\chisq$ values. The favoured value for $\\fvn$ is that for which the $\\chisq$ is minimized. ",
        "conclusions": "\\label{discussion} Some issues with the use of GRBs as standard candles for constraining cosmological parameters are discussed by \\citet{ggf2006}, \\citet{g2009}, and \\citet{p2009}. In particular, the correlations between the isotropic luminosities and the luminosity indicators are weak in a $\\chisq$ sense---physical factors unaccounted for are producing large scatter. Strictly speaking, the $\\Epeak-\\Egamma$ correlation is the only relationship with a sufficiently low reduced $\\chisq$ to admit cosmological parameter estimation, albeit with the caveat that the measurement of the jet break time assumes a particular fireball model~\\citep{g2009}. \\citet{p2009} point out that the luminosity correlations are statistical in nature, rather than being, as they should ideally, one-to-one relations between uncorrelated quantities. This meant, for example, that the comparatively tight $\\Epeak-\\Egamma$ correlation found by \\citet{g2004} actually weakened with the introduction of more data points.  Systematic uncertainties such as dust extinction and evolution constitute considerable limitations to cosmological parameter estimation with \\sne. Many of these uncertainties, for example Malmquist bias and gravitational lensing, are considered negligible in the redshift range over which \\sne\\ occur. For the redshift range over which GRBs occur, one would expect that Malmquist bias and lensing might cause at least some of the scatter in the GRB Hubble diagram, but these biases are shown in S07 to be negligibly small. Obscuration by dust is not an issue for GRBs~\\citep{g2004}, but selection and evolution effects can potentially influence the current GRB sample. The well-known ``Amati'' correlation for long-duration GRBs between isotropic-equivalent radiated energy $\\Eiso$, describing the intensity of the burst, and the photon energy at which the time-averaged spectrum peaks $\\Epi$, although proving to be quite robust~\\citep*{amati08,amati10}, has shown evidence of variation with redshift~\\citep{li07} and susceptibility to detector threshold selection effects~\\citep{butler2007}. \\citet{p2009} find evidence for evolution of the GRB peak luminosities, but this should not affect the Hubble diagram, since it is the luminosity relations which should give the right distances for placement on the Hubble diagram.  It can be argued that the kind of relativistic and geometric effects that underlie the luminosity relations should not be greatly affected by evolution or the metallicity of the progenitor~\\citep{s07}. It is conceivable that a better understanding of GRB physics in the future will allow them to be used as ``standardizable'' candles, and put their utility for cosmological parameter estimation and discrimination between cosmological models on a firmer basis. Ongoing observational programmes such as \\emph{Swift} continue to contribute to this aim. We need to know more about the physics of the GRBs, and we need more high-quality measurements of GRB redshifts, light curves and spectra.  In the meantime, however, we can obtain glimpses of the potential applications of standard candles whose range extends into the era of decelerating cosmic expansion. In the present study, the correlations are forced to be a good fit by incorporating the additional ``systematic error'' term, computed such that it makes the $\\chisq$ of the best-fit correlation equal to one. This term contributes (in quadrature) to the uncertainty in the log of the isotropic luminosity which propagates through to the Hubble diagram $\\chisq$ via eqs (\\ref{eq:pbolounc}) and (\\ref{eq:sbolounc}). A single GRB at a redshift of 5 or 6, with better-determined physical characteristics, potentially carries more statistical power than a single \\sn\\ at $z=1.7$ because of the Hubble diagram ``lever arm''---the Hubble diagram at redshifts $z>2$ changes with a different cosmological model or cosmological parameters much more than it does at lower redshifts. We obtain results that are not inconsistent with current models, with certain acknowledged caveats. In particular, there is much scope for progress in improving the GRB Hubble diagram, and much to be gained."
    },
    "1107/1107.2463_arXiv.txt": {
        "abstract": "We observed at 22~GHz with the VLBI array VERA a sample of 1536 sources with  correlated flux densities brighter than 200~mJy at 8~GHz. One half of target sources has been detected. The detection limit was around 200~mJy. We derived the correlated flux densities of 877 detected sources in three ranges of projected baseline lengths. The objective of these observations was to determine the suitability of given sources as phase calibrators for dual-beam and phase-referencing observations at high frequencies. Preliminary results indicate that the number of compact extragalactic sources at 22~GHz brighter than a given correlated flux density level is twice less than at 8~GHz. ",
        "introduction": "Currently, very long baseline interferometry (VLBI) astrometry is the best tool to measure distances and motions of sources located at kpc scale and hence, to explore the structure of the Milky Way in the Galactic scale. For instance, Japanese VERA project (VLBI Exploration of Radio Astrometry; \\citet{r:vera00}) has been conducting astrometric monitoring of positions of Galactic maser sources with respect to reference compact extragalactic objects, yielding handful measurements of parallaxes and proper motions of maser sources (e.g., see recent PASJ special issue for VERA, such as \\citet{r:hon11,r:nag11} and others). The Very Long Baseline Array (VLBA) is actively used for astrometry of Galactic maser sources (\\citep[e.g.,][]{r:reid09} and the currently ongoing Bar and Spiral Structure Legacy (BeSSeL) survey) and the European VLBI Network (EVN) conducts astrometric observations of methanol maser sources \\citep[e.g.,][]{r:ryg10a}.  In order to measure parallax and proper motion of a radio source at kpc scales, it is observed in the phase-referencing mode by frequent switching pointing between the target and a calibrator source. This technique significantly reduces phase variations caused by tropospheric fluctuations. To do this effectively, calibrators must be located close to target sources, typically within 1--2\\degr\\ separation. This requires a high density of calibrator sources in the sky, and hence, there is still a strong demand for finding many calibrator sources.  To date, there have been several massive surveys of compact calibrators such as VCS (VLBA Calibrator Surveys), \\citet{r:vcs6} and references therein), the LCS (Long Baseline Array Calibrator Survey) for the southern hemisphere \\citep{r:lcs1}, VIPS (VLBA Image and Polarization Survey) \\citep{r:vips,r:astro_vips}, and several ongoing programs: the program of study the {\\it Fermi} active galaxy nuclea (AGNs) at parsec scales{\\footnote{\\web{http://astrogeo.org/faps}}} (Kovalev et al. (2011), paper in preparation), the program of observing radio-loud 2MASS (Two Micron All Sky Survey) galaxies{\\footnote{\\web{http://astrogeo.org/v2m}}} \\citep{r:con11}, the program of observing optically bright quasars \\citep{r:bou08,r:bou11,r:obrs1}, and the recent VLBA calibrator search for the BeSSeL survey \\citep{r:br145a}.  Together with regular geodetic VLBA observations of 1000 sources \\citet[the RDV program][]{r:rdv}), by June 2011 positions of 6455 sources at a milliarcsecond level of accuracy were derived from analysis of these massive surveys. The sources turned out compact enough to be detected with VLBI, i.e. they have a core of mas scale. However, these surveys were in most cases conducted in relatively low frequencies such as 2 (S-band), 5 (C-band) or 8~GHz (X-band), at which the telescope performance is the best. On the other hand, recent VLBI maser astrometry is often done at frequencies higher than 10~GHz. For instance, VERA's main bands are 22 (K-band) and 43~GHz (Q-band) for H$_2$O and SiO maser sources. Maser astrometry with VLBA is mainly conducted at 12~GHz for CH$_3$OH masers and 22~GHz for H$_2$O masers. Therefore, calibrator information at high frequencies (such as K and higher bands) is of great importance for on-going and future astrometric observations. Compact calibrators which are cores of radio bright AGNs have a wide variety of their spectra: for the majority of sources the correlated flux density decreases with the frequency, although some sources may have spectra growing with frequency or peaking within the GHz regime. Hence, the extrapolation of the correlated flux densities from S and X band to 22 or 43~GHz is highly unreliable. For successful phase-reference or dual-beam observations, the correlated flux density should be known with accuracy at least 30\\% in order to correctly predict the signal-to-noise ratio (\\SNR). Therefore, it is necessary to conduct a systematic survey of K-band flux densities for the compact sources which were already detected in S and X bands.  We have identified $\\sim\\!\\!2000$ sources previously observed with VLBI with $\\delta > -30\\degr$ with correlated flux densities $>200$~mJy at X-band at baselines longer than 900~km. Analysis of the dependence of the number of sources $N$ with the correlated flux density exceeding $S$ as a function of $S$ suggests that this sample is complete at the 95\\% level (Kovalev 2010, private communication). Of these sources, 511 have been previously observed in large K-band surveys: VERA Fringe Search Survey \\citep{r:vera_fss}, KQ survey \\citep{r:kq}, VLBI Galactic plane survey (VGaPS) \\citep{r:vgaps}, in the EVN Galactic plane survey (EGaPS) \\citep{r:egaps}, and their correlated flux densities at 22~GHz have been measured. The K-band brightness of other objects was not known.  We conducted a dedicated survey of remaining 1536 sources at 22~GHz with VERA in the K-band Calibrator Survey (KCAL) campaign. The goal of these observations was to check their detectability at K-band and to measure the correlated flux densities of detected sources at baselines 1000--2000~km.  The first objective of this campaign was to provide a complete list of calibrators suitable for VERA observations of faint targets. According to our prior observations, the detection limit of the VERA network for 2 minutes of integration time is around 200~mJy, depending on weather conditions. Therefore, the list of sources observed in this and the previous K-band surveys is expected to approach the completeness at the 200~mJy level, provided the spectra of compact cores are flat or falling. According to \\citet{r:mass10} who analyzed simultaneous ATCA spectra at 4.8, 8.4 and 20~GHz, the share of sources with growing spectra that may be missed in our sample does not exceeded 8\\%.  The second objective of this campaign is to collect information for a population analysis of a large complete sample. In particular, the analysis of the dataset that combines existing and new data will help to answer the question what is the distribution of spectral indexes of the core regions and the source compactness at high frequencies, whether the spectral index at parsec scales is systematically different than the spectral index at kiloparsec scales, and whether the compactness at K-band is systematically different than the compactness at X and S bands.  In this paper we present results of the survey. In section \\ref{s:obs} we describe the observations, their design and scheduling. In section \\ref{s:anal} we discuss analysis technique. The catalogue of correlated flux densities of detected sources accompanied with analysis of flux density uncertainties is presented in section \\ref{s:cat} followed by concluding remarks that are given in section \\ref{s:sum}. ",
        "conclusions": "\\label{s:sum}  We observed with VERA at 22~GHz a subset of the complete sample of continuum compact extragalactic sources with correlated flux densities $> 200$~mJy at X-band at declinations $>-30\\degr$. The subset excluded the sources previously detected at K-band at large VLBA and VERA surveys. Of 1536 target sources, approximately one half has been detected. The errors of the correlated flux densities are a level of 20\\%.  \\begin{figure}[h] \\caption{\\ifpre \\rm \\fi The distribution of the correlated flux densities at baseline projection lengths longer than 100 megawavelengths. The last bin of the histogram has all the sources with correlated flux density $>2$ Jy. } \\label{f:distr} \\par\\medskip\\par \\ifdraft \\includegraphics[width=0.46\\textwidth]{kcal_distr_200.eps} \\else \\includegraphics[width=0.46\\textwidth]{kcal_distr.eps} \\fi \\end{figure}  Figure~\\ref{f:distr} shows the distribution of the KCAL correlated flux densities. Assuming that the parent population of sources is uniform in the range of flux densities 1--1000~mJy, we explain the sudden drop in the number of sources with correlated flux densities below 200~mJy as an indication of under-representation of sources weaker than that limit in the catalogue because they are not reliably detected with VERA. Thus, the KCAL is incomplete at flux densities below 200~mJy. This result is in agreement with our analysis of previous VERA observations \\citep{r:vera_fss} where we estimated the probability of detection of a source with the correlated flux density 200~mJy at a level of 70\\%.  \\begin{figure}[h] \\caption{\\ifpre \\rm \\fi The distribution of the spectral indices $\\alpha$ ($F(\\nu) \\sim \\,v^{\\alpha}$) of 1536 sources from the input catalogue. The spectral index was computed from median correlated flux densities at 8.6 and 2.3~GHz at baseline projected lengths shorter than 900~km. } \\label{f:kspi} \\par\\medskip\\par \\ifdraft \\includegraphics[width=0.47\\textwidth]{kcal_spind_distr_200.eps} \\else \\includegraphics[width=0.47\\textwidth]{kcal_spind_distr.eps} \\fi \\end{figure}   The detection limit of VERA at 22~GHz, 200~mJy, corresponded the lowest correlated flux density of the input source list, 200~mJy, {\\it at 8~GHz}. The majority of the sources from the input list were previously observed at VLBI at both 8.6 and 2.3~GHz. The distribution of spectral indices ($F(\\nu) \\sim \\,v^{\\alpha}$) of the compact component of these sources shows a peak near spectral index 0 (see Figure~\\ref{f:kspi}). Among 1536 sources from the input list, 48\\% had spectral index greater then zero, and therefore, their extrapolated flux density at 22~GHz was greater than 200~mJy, the average detection limit of the KCAL survey. Although  a measured correlated flux density at 22~GHz for an individual source may be less of greater than the flux density extrapolated from 8.6/2.3~GHz, if to consider the entire population as a whole, the measured flux density at 22~GHz turned out {\\it on average} very close to the extrapolated one.  Results of the KCAL survey augmented with results of prior K-band surveys form the list of objects with known correlated flux densities. By June 2011, this list\\footnote{Available at \\web{http://astrogeo.org/rfc}} contained 1161 objects. Among these sources, 766 objects have correlated flux densities greater than 200~mJy at baselines shorter than 70~M$\\lambda$ and 608 objects are brighter than 200~mJy at baselines longer than 100~M$\\lambda$. These sources are considered as a pool of calibrators for VERA in 2011. After completion of a planned sensitivity upgrade, even weaker sources can be used for calibrators.  We reserve a rigorous population analysis to a future publication. Preliminary results indicate that the number of compact extragalactic sources at K-band brighter than a given correlated flux density level is twice less than at the X-band.   We would like to thank Alan Fey for making publicly available not only contour plots of images from the KQ survey, but brightness distribution and calibrated flux densities in the FITS-format. The availability of this information was crucial for our project."
    },
    "1107/1107.2180_arXiv.txt": {
        "abstract": "The 4.9 GHz Micro-Arcsecond Scintillation-Induced Variability (MASIV) Survey detected a drop in Interstellar Scintillation (ISS) for sources at redshifts $z \\gtrsim 2$, indicating an apparent increase in angular diameter or a decrease in flux density of the most compact components of these sources, relative to their extended emission. This can result from intrinsic source size effects or scatter broadening in the Intergalactic Medium (IGM), in excess of the expected $(1+z)^{1/2}$ angular diameter scaling of brightness temperature limited sources resulting from cosmological expansion. We report here 4.9 GHz and 8.4 GHz observations and data analysis for a sample of 140 compact, flat-spectrum sources which may allow us to determine the origin of this angular diameter-redshift relation by exploiting their different wavelength dependences. In addition to using ISS as a cosmological probe, the observations provide additional insight into source morphologies and the characteristics of ISS. As in the MASIV Survey, the variability of the sources is found to be significantly correlated with line-of-sight H$\\alpha$ intensities, confirming its link with ISS. For 25 sources, time delays of about 0.15 to 3 days are observed between the scintillation patterns at both frequencies, interpreted as being caused by a shift in core positions when probed at different optical depths. Significant correlation is found between ISS amplitudes and source spectral index; in particular, a large drop in ISS amplitudes is observed at $\\alpha < -0.4$ confirming that steep spectrum sources scintillate less. We detect a weakened redshift dependence of ISS at 8.4 GHz over that at 4.9 GHz, with the mean variance at 4-day timescales reduced by a factor of 1.8 in the $z > 2$ sources relative to the $z < 2$ sources, as opposed to the factor of 3 decrease observed at 4.9 GHz. This suggests scatter broadening in the IGM, but the interpretation is complicated by subtle selection effects that will be explored further in a follow-up paper. ",
        "introduction": "It is well established that the intraday variability (IDV) observed in many compact, flat-spectrum active galactic nuclei (AGN) at centimetre wavelengths is predominantly caused by scintillation in the turbulent and ionized interstellar medium (ISM) of our own Galaxy. The idea was first proposed by \\citet{heeschenrickett87} in order to resolve the brightness temperature problem in AGN, where intrinsic variability on the time-scales observed implied brightness temperatures well over the $10^{12}$ K inverse-Compton limit for incoherent synchrotron emission \\citep{hunstead72,heeschen84}. Substantial observational evidence has accumulated in the last decade to support this theory. Time delays of up to 8 minutes have been observed in the scintillation patterns of the most rapid scintillators at widely spaced telescopes \\citep{jaunceyetal00, dennett-thorpedebruyn02, bignalletal06}, as would be expected of interference patterns drifting across the surface of the Earth as a result of relative motion between the ISM and the Earth. Annual cycles have also been detected in AGN variability time-scales \\citep{rickettetal01, jaunceymacquart01, bignalletal03, dennett-thorpedebruyn03, jaunceyetal03}, interpreted as being modulated by the orbital motion of the Earth around the Sun. When the Earth's motion is parallel to the motion of the scattering medium, the variability time-scales are longer, while shorter timescale variability occurs when the Earth's motion is anti-parallel to that of the scattering medium.   More recently, the Microarcsecond Scintillation-Induced Variability (MASIV) Survey \\citep{lovelletal03, lovelletal08} provided further confirmation of interstellar scintillation (ISS) as the principal mechanism behind IDV. In a 4.9 GHz survey of more than 500 compact, flat-spectrum sources at 4 separate epochs spaced throughout a year, it was found that more than half of the sources were variable in at least one epoch. The survey showed a strong correlation between AGN variability and Galactic latitudes, as well as line of sight H$\\alpha$ intensity (as an estimate of the emission measure of the ISM), thus strengthening the link between IDV phenomena and the ionized ISM.   The effects of scintillation are highly dependent on the angular size of the source. Compact sources tend to scintillate more than extended sources, analogous to the fact that `stars twinkle, but planets do not' when observed through the Earth's atmosphere. For weak interstellar scattering (WISS), which is usually the case for scintillating sources at frequencies $\\geq$ 5 GHz at mid-Galactic latitudes, the angular size of the source must be comparable to or smaller than the size of the first Fresnel zone \\citep{narayan92} at the scattering screen. The variations in the flux density of the MASIV sources indicate that a significant portion of the emission comes from compact components with angular diameters on a scale of 10 to 50 microarseconds. Follow-up observations of the morphologies of scintillating and non-scintillating sources using Very Long Baseline Interferometry (VLBI) confirm that, at milliarcsecond scales, the scintillating sources are more core-dominated than non-scintillating sources \\citep{ojhaetal04a, ojhaetal04b}, where a large proportion of their flux densities are confined in an ultra-compact core region, best interpreted as highly Doppler-boosted jet emission with intrinsic brightness temperatures close to the inverse-Compton limit of $10^{12}$ K. In the presence of stronger, milliarcsecond-scale jets, scintillation effects are diminished.   A most tantalizing result of the MASIV survey was the discovery of a significant drop in the fraction of scintillating sources, as well as their variability amplitudes, at redshifts above 2 \\citep{lovelletal08}. While the angular diameters of a population of sources limited by a maximum brightness temperature are expected to scale with $(1+z)^{1/2}$ in an expanding $\\Lambda$CDM Universe \\citep{rickettetal07}, due to the source brightness temperature in the observer's frame being a factor of $(1+z)$ lower relative to the source brightness temperature in the comoving frame, the redshift dependence observed in the survey was found to be in excess of this effect. Therefore, it can be attributed either to an additional increase in typical source angular diameters, or a decrease in flux densities of the aforementioned ultra-compact components of the AGN relative to their more extended components. An increase in the angular diameter can be a result of angular broadening as the radio waves propagate through the turbulent, ionized intergalactic medium (IGM), or a decrease in the Doppler boosting factor in AGN jets at earlier epochs. A reduction in flux density of the ultra-compact component can also be due to a decrease in the Doppler boosting factor of the jets or even a decrease in the prevalence of such very compact objects at higher redshifts.   Identifying with certainty the cause of this redshift dependence in AGN ISS has profound cosmological implications. If the cause is intrinsic to the sources, it provides new results on the evolution of AGN morphologies at scales two orders of magnitude finer than that available to VLBI. On the other hand, detection of angular broadening in the IGM would provide a new observational tool for the study of the majority of baryons (90\\% are thought to reside in the IGM, see e.g. \\citet{fukugitapeebles04}) and their evolution after the epoch of reionization. Being sensitive to turbulent, ionized components of the IGM, it will thus complement Ly$\\alpha$ studies of the IGM which are sensitive only to the neutral component, as well as UV and X-ray observations of the hot intra-cluster medium (ICM). Even more exciting is the prospect of detecting the elusive warm-hot intergalactic medium (WHIM), predicted by cosmological hydrodynamical simulations \\citep{cenostriker99, daveetal01, cenostriker06} to exist in an almost fully ionized state at temperatures of $10^{5}$ to $10^{7}$ K due to gravitational shock heating. The WHIM  has so far been very difficult to detect conclusively \\citep{bregman07}, and forms one of the key science drivers of the next generation of X-ray instruments.   To determine the origin of this redshift dependence in ISS, we have conducted multi-frequency observations of a sub-sample of the MASIV sources using the Very Large Array (VLA). Observing at multiple frequencies can potentially provide an effective technique for discriminating the cause of the redshift dependence in AGN scintillation from among the three possible explanations - cosmological expansion, angular broadening in the IGM and evolution of AGN jets. Scatter broadening in the ionized IGM should have a stronger wavelength dependence compared to the angular size-wavelength dependence of the source core, leading to a decrease in the redshift scaling at higher frequencies. These observations, therefore, have the ultimate goal of using ISS as a cosmological probe - potentially of AGN jet evolution, turbulence in the IGM, or the curvature of the Universe.   While determining the cause of the redshift dependence of ISS remains the main objective of the study, the opportunity to gain additional insight into ISS phenomena and AGN morphology provides additional motivations for conducting the observations. Firstly, the 11-day duration of these observations (as opposed to the 3 or 4 day epochs in the MASIV survey) gives improved constraints on source scintillation timescales. Secondly, multi-frequency observations of ISS provide a means of detecting any angular offset in the positions of the AGN cores at different frequencies, observed as a delay in the scintillation patterns between frequencies as the scattering screen drifts across the source, as seen in PKS 1257--326 \\citep{bignalletal03}. Such frequency core-shifts have also been observed in VLBI studies \\citep{kovalevetal08}, and are interpreted in terms of opacity effects in the source jet. Thirdly, multi-frequency observations also enable us to estimate the instantaneous spectral indices of the sources based on concurrent mean flux densities to study its relationship to ISS (further details in Section~\\ref{spectralindexsec}). Finally, together with the MASIV Survey, the experiment enables us to place a lower limit on the detectability of ISS amongst the presence of noise and other systematic errors using the VLA, in addition to providing a platform for exploring various methods of estimating and accounting for these errors. These observations thus act as a demonstrator for future large scale surveys such as the planned Australian Square Kilometre Array Pathfinder (ASKAP) Variables and Slow Transients (VAST) Survey \\citep{murphyandchatterjee09}.   Section~\\ref{observations} of this paper describes the observations and data reduction process for this follow-up to MASIV. This is followed by a detailed elucidation of the various methods used in the analysis of the data (including error estimation and correction), along with the results (Section~\\ref{analysisresults}). Section~\\ref{conclusion} presents our conclusions. Further interpretation of these results with regards to the redshift dependence of ISS in AGN and its cosmological implications is discussed in a separate paper, which will investigate the possible source selection effects that may lead to biases in the interpretation. ",
        "conclusions": "Multi-frequency observations of 140 flat spectrum AGN were carried out using the VLA for a total duration of 11 days. The sensitivity of the VLA and careful calibration enabled noise and systematic errors to be reduced down to a level of $\\lesssim 1\\%$. These errors were then quantified as a quadratic sum of the noise ($s$) and calibration errors ($p$), and subsequently subtracted from the SF values obtained by fitting a simple model to the SFs calculated from the source lightcurves. Statistically significant correlations with H$\\alpha$ intensities were obtained for the SFs at all time-lags (using both model and single sample estimates calculated from the lightcurves) and at both frequencies, confirming the MASIV results linking the variability of the sources to ISS. Cross-covariance functions between source lightcurves at 4.9 GHz and 8.4 GHz reveal that the patterns of scintillating sources are correlated. A time-lag shift in the peak of the cross-covariance function points to the possibility of a core shift in such sources at different frequencies. Although there were no clear trends with regards to mean spectral indices above -0.4, a clear drop in ISS amplitudes was observed for sources with spectral indices below -0.4, confirming reduced scintillation in steeper spectrum sources. As in the MASIV survey, a drop in ISS at $z \\gtrsim 2$ was observed at 4.9 GHz. Of even greater significance is the detected reduction in the redshift dependence of ISS at 8.4 GHz, suggestive of scatter broadening in the IGM if weak ISS is assumed. A follow-up paper will delve further into its interpretation, pending a full investigation into source selection effects and comparisons with more accurate models to understand the complex underlying physics.   The results of MASIV and this study continue to demonstrate the potential of using ISS as an astrophysical and cosmological probe. As shown, ISS can be used to estimate the core-shifts of radio sources at a higher resolution than that of VLBI. Such observations will be important for astrometric applications in the selection of sources for an International Celestial Reference Frame. Combining multi-frequency scintillation observations with VLBI imaging (to obtain the true angular separation between the cores) raises the prospect of putting constraints on actual scattering screen distances and velocities, providing further insight into the physics of the ISM.   While this study included only 140 sources, future large scale surveys of IDV are already being planned, such as the VAST Survey \\citep{murphyandchatterjee09} - one of the key survey science projects of the ASKAP. These future experiments using the ASKAP will operate at a much higher survey speed. Thus the various techniques applied in this study, in our effort to obtain the best possible characterization of the variability of the sources, provide valuable insight for these future surveys, which require the development of efficient pipelined algorithms for the calibration, detection and analysis of the large quantities of data."
    },
    "1107/1107.3186_arXiv.txt": {
        "abstract": "\\noindent A diagrammatic approach to calculate $n$-point correlators of the primordial curvature perturbation $\\zeta$ was developed a few years ago following the spirit of the Feynman rules in Quantum Field Theory.  The methodology is very useful and time-saving, as it is for the case of the Feynman rules in the particle physics context, but, unfortunately, is not very well known by the cosmology community.  In the present work, we extend such an approach in order to include not only scalar field perturbations as the generators of $\\zeta$, but also vector field perturbations.  The purpose is twofold:  first, we would like the diagrammatic approach (which we would call the Feynman-like rules) to become widespread among the cosmology community;  second, we intend to give an easy tool to formulate any correlator of $\\zeta$ for those cases that involve vector field perturbations and that, therefore, may generate prolonged stages of anisotropic expansion and/or important levels of statistical anisotropy.  Indeed, the usual way of formulating such correlators, using the Wick's theorem, may become very clutter and time-consuming. ",
        "introduction": "% Our Universe exhibits departures from the exact isotropy and homogeneity;  we may see it from the distribution of temperatures in the cosmic microwave background radiation (CMB) \\cite{wmap7} and from the distribution of matter that gives rise to the large-scale structure \\cite{sdss}.  The origin of such departures is a marvelous phenomenon that involves two outstanding mechanisms in particle cosmology \\cite{mukhanovbook}:  the production of virtual particles via the Heisenberg uncertainty principle, and the primordial accelerated expansion of the Universe that defines a dynamical particle horizon.  The joint work of these two mechanisms leads to the production of real particles (the Hawking radiation phenomenon) and the classicalisation of the field perturbations that live in the spacetime which describes our Universe. The nature of such field perturbations may be scalar, vectorial, spinorial, or even, it may correspond to p-forms\\footnote{The particle production process has been well studied in the scalar field case \\cite{mukhanovbook}.  Some nice studies in the vector field case may be found in Refs. \\cite{dklr,dimovector,dimosupervector,dkw}.}.  In addition, their contribution to the departure from the exact isotropy and homogeneity (parameterized by the primordial curvature perturbation $\\zeta$ \\cite{weinbergbook,lythbook}) may be operative during inflation (the standard inflationary mechanism \\cite{mukhanovchisov,hawking,starobinsky,guthpi,bardeen,fischler}), at the end of inflation (the inhomogeneous reheating mechanism \\cite{dvali}), or after inflation (the curvaton mechanism \\cite{lythwands,moroi,lythwands2}).  Depending on the nature of the field perturbations, on the number of fields involved, and on the action describing the theory of gravity and the matter content of the Universe, the statistical properties of the distribution of temperatures in the CMB and of the distribution of matter that form the large-scale structure may exhibit different features: once the variance has been fixed, the power spectrum of $\\zeta$ may become scale-dependent, either blue-tilted or red-tilted, or even may become dependent on the direction of the wavevector (signaling violation of the rotational invariance in the two-point correlator of $\\zeta$, a property which is called statistical anisotropy \\cite{dklr,abramo}).  The higher-order correlators of $\\zeta$ are also affected:  the odd-point correlators may become different to zero (signaling departures from a gaussian distribution \\cite{lythbook}) and the even-point correlators may become different to products of two-point correlators (signaling also non-gaussianity \\cite{lythbook}).  Moreover, the three-point correlator may depend on the directions of the three wavevectors, something which has been called anisotropic non-gaussianity \\cite{dimopoulos, Yokoyama:2008xw}. Talking a bit more about the statistical anisotropy, it is generated by either a primordial anisotropic expansion, or a non-scalar nature of the fields involved, or both.  Calculating the $n$-point correlators of $\\zeta$ from a well-defined action is crucial in order to compare theory and observation\\footnote{At the end, the theory cannot predict the exact outcome of an experiment in a member of the ensemble since the quantum nature only allows us to predict probabilities in the ensemble.  The ergodic theorem is crucial in this respect \\cite{weinbergbook} and, therefore, the assumption of statistical homogeneity (translational invariance of the $n$-point correlators) must be maintained.}. The right framework to propagate the statistical properties of the field perturbations to the statistical properties of $\\zeta$ is the cosmological perturbation theory (CPT) \\cite{weinbergbook,kodamasasaki,mukhanovreport,malikreport}; however, this normally involves lengthy calculations, even more if the nature of the fields is not scalar \\cite{himmetoglu4,watanabe2,dulaney,peloso}.  In addition, interesting phenomena such as non-gaussianity may only be accessible if the CPT is taken to second or higher orders \\cite{lythbook,lythyeinzon1,yeinzonbook}. Despite of this, the CPT is valid throughout all scales, leaving no room for discrepancies attributed to not considered subhorizon phenomena.  A different approach is via the $\\delta N$ formalism \\cite{starobinsky2,ss,tanaka,sasaki1}, where $\\zeta$ is identified with the perturbation in the number of e-folds of expansion from an initial time in a flat slicing to a final time in a uniform energy density slicing (the threading must be comoving).  The $\\delta N$ formalism gives an expression for $\\zeta$ which is valid to all orders in CPT;  however, it is only valid for superhorizon scales (in absolut contrast with CPT) \\cite{lythbook}.  Of course, extracting the statistical properties of the distribution of $\\zeta$ in the $\\delta N$ formalism requires also to do some ``perturbation theory'': to expand $\\zeta \\equiv \\delta N$ in a Taylor series and to cut it out at the desired order \\cite{yeinzonbook,lythyeinzon2}.  The classicalisation is given at superhorizon scales \\cite{lythseery}, as well as the conservation of $\\zeta$ is if the adiabatic pressure condition is hold \\cite{lythbook,sasaki1};  that is why we do not worry much about the subhorizon scales and still we obtain very precise results at a very cheap computational cost.  The way of calculating $n$-point correlators of $\\zeta$ in the $\\delta N$ formalism is very similar to the way of calculating scattering amplitudes in the cannonical formulation of Quantum Field Theory \\cite{peskin}:  the Wick's theorem is essential \\cite{wick}, the calculation is very direct but not so intuitive, and anyway it may become very clutter and time-consuming. Feynman was very clever, even for minor aspects such as avoiding not so intuitive procedures when calculating things;  he developed a series of diagrammatic rules \\cite{feynman1,feynman2}\\footnote{Actually, this technique was first described by Feynman at the Poconos Conference in 1948, and published in 1951 \\cite{feynman3}.} that can reproduce every piece of the calculation of scattering amplitudes, very easy to apply, very easy to remember, and very illuminating at trying to visualize what is actually happening in a quantum process.  Feynman rules became widespread among the particle physicists community and it is nowadays the standard language to interpret quantum processes and calculate scattering amplitudes \\cite{diagrammatica}. In view of the above, a few years ago a diagrammatic approach to calculate $n$-point correlators of $\\zeta$ was developed including only scalar fields as the generators of $\\zeta$ \\cite{Byrnes:2007tm}\\footnote{A closely related diagrammatic formalism was developed some time ago in order to study large-scale structure formation via gravitational instability; see for instance Ref. \\cite{Crocce:2005xy} and references therein. See also Ref. \\cite{Giddings:2010ui} for a diagrammatic treatment of perturbative calculations in the ``in-in\" formalism and Refs. \\cite{Ensslin:2008iu,Ensslin:2010bw} for a diagrammatic approach to information field theory applied to the reconstructions of non-linear signals.}.  Such ``Feynman-like'' rules share the same advantages of its counterpart in Quantum Field Theory, making very easy and intuitive to start any $n$-point correlator calculation.  External lines in this case are identified as Fourier modes of $\\zeta$ and internal lines, the propagators, are identified as the power spectra of the scalar field perturbations.  If the field perturbations are non-gaussian, new diagrams appear which do not have a correspondence in the particle physics Feynman rules:  internal lines that split like branches of a tree;  they correspond to connected $n$-point correlators of the field perturbations.  The issue of the vertex renormalization was also treated, showing in an elegant way how to absorb the diagrams with dressed vertices into just one diagram with undressed, but redefined, vertices, in complete analogy with the Feynman rules case \\cite{diagrammatica}\\footnote{The standard vertex corresponds to a $m$-order derivative (being $m$ the number of internal lines ending in that vertex) of the unperturbed number of e-folds $N$ with respect to the fields involved \\cite{Byrnes:2007tm}.}.  In this paper, we want to draw the attention of the reader into two significant aspects of the history of the Feynman-like diagrammatic approach in cosmology.  First, although very powerful as it is in the particle physics context, not many cosmologists are aware of the existence or the power of the Feynman-like rules, and still most of the calculations based in the $\\delta N$ formalism are performed directly, invoking the Wick's theorem.  Of course, the arrival to the relevant integrals is lengthy, the intuitive connection between an important feature in the calculation and the physics behind it is almost lost, and the scene quickly becomes very clutter.  Second, vector field perturbations were not considered in the original formulation of the Feynman-like rules.  Vector fields as possible generators of $\\zeta$ have attracted the attention of several scientists as they may generate significant statistical anisotropy, still in agreement with present observations \\cite{groeneboom,hanson,hanson2,ma}, via the anisotropic expansion they may generate \\cite{himmetoglu4,watanabe2,dulaney,peloso,watanabe1} and also via its own vector nature \\cite{dklr,dimovector,dimosupervector,dkw,jabbari1,jabbari2,kostascurvaton2,kostasattractive,karciauskas}.  Indeed, statistical anisotropy surely will become discriminator of models for the generation of $\\zeta$ taking into account the forthcoming increment in the precision of observations starting with the PLANCK satellite data \\cite{abramo,pullen}.  The first purpose of this paper is to try to disseminate the diagrammatic approach to cosmologists as much as possible; time will tell if this objective was reached.  The second purpose of this paper is to extend the Feynman-like rules to the case where the contributions to $\\zeta$ are due to several scalar and vector fields, producing the latter, in general, anisotropic expansion, statistical anisotropy, and anisotropic non-gaussianity \\cite{dimopoulos,vrl,vr,mythesis,bartolo1,bartolo2,bartolo3}.  We try to be as general as possible, allowing for an arbitrary number of fields, taking into account the transverse as well as the longitudinal polarizations of the vector fields (i.e. allowing for mass terms), keeping in mind that the action may not be parity-invariant, considering, in general, anisotropic expansion, and assuming that the probability distribution functions for the scalar field perturbations, as well as for the vector field perturbations\\footnote{Actually, the probability distribution functions for the scalar perturbations that multiply the respective polarization vectors.}, may in general be non-gaussian.  It is worth emphasizing that, in the context of nowadays cosmology, it is very important to have an efficient method for the calculation of loop corrections; this is because high precision cosmological measurements could be able to access observational signatures encoded in the loop diagrams if they turn to be non-negligible \\cite{Sugiyama:2011jt}, so it is crucial to go beyond tree level terms. Moreover, from the theoretical point of view, there exist scenarios in which loop contributions are larger than the tree-level terms and, therefore, are able to produce large non-gaussianities \\cite{vrl, vr, Boubekeur:2005fj, Cogollo:2008bi, Rodriguez:2008hy, Kumar:2009ge}; in those scenarios, it is essential to have an efficient method to evaluate  such contributions.  The layout of the paper is as follows.  In Section 2 the $n$-point correlators are introduced;  the meaning of statistical homogeneity, statistical isotropy, and gaussianity is clearly established.  In Section 3 the $\\delta N$ formalism is introduced, allowing for the possibility of anisotropic expansion and, therefore, for the existence of vector fields as generators of $\\zeta$;  the two-point correlator of $\\zeta$ is calculated using the non-diagrammatic technique (i.e. by employing the Wick's theorem) up to the one-loop level.  Section 4 is devoted to the presentation of the Feynman-like rules in the limit where the field perturbations are gaussian; the calculation of the two- and three-point correlator of $\\zeta$ up to the one-loop level, and four-point correlator up to the tree level, are performed employing the diagrammatic approach and compared with the findings in Section 3.  Section 5 introduces the rules in the non-gaussian case; again, some examples are considered and compared with the non-diagrammatic approach.  In Section 6 the vertex renormalisation is studied for the gaussian and non-gaussian cases.  The conclusions of this work are presented in Section 7. ",
        "conclusions": "% Formulating the $n$-point correlators of the primordial curvature perturbation $\\zeta$ in terms of integrals involving correlators of the field perturbations is a very direct procedure; however it is in general very time-consuming, clutter, and not so intuitive.  The same situation happened with the calculation of scattering amplitudes in cannonical Quantum Field Theory until Feynman exposed for the first time his diagrammatic rules \\cite{feynman1,feynman2,feynman3} which, since then, have become the standard way of doing such a kind of calculations \\cite{diagrammatica}; the reason: Feynman rules are vey easy to remember, very easy to implement, very intuitive, and allow us to understand the physics behind a quantum process.  A few years ago, the ``Feynman-like rules'' to calculate $n$-point correlators of $\\zeta$ were developed \\cite{Byrnes:2007tm}, with the same aims as the Feynman rules in the particle physics context.  In that work, only scalar fields were introduced as the generators of $\\zeta$.  Since then, not many cosmologists have used this diagrammatic approach, wasting the interesting properties stated above.  As mentioned in the introduction, the relevance of having an efficient method to calculate loop corrections relies mainly in the possibility that observational data coming from high precision cosmological probes could open a window to access effects related to such corrections; there are also scenarios in which evaluating loop contributions is essential because those terms dominate over the tree level ones and constitute a source for large and observable non-gaussianities \\cite{vrl, vr, Boubekeur:2005fj, Cogollo:2008bi, Rodriguez:2008hy, Kumar:2009ge}.  In the present work, we have been interested in promoting even more this methodology among the cosmology community. In addition, we have extended it to the interesting case where multiple scalar and vector fields contribute to the generation of $\\zeta$; in this situation it is possible to obtain prolonged stages of anisotropic expansion (see for instance Ref. \\cite{watanabe1,anexp1,anexp2,anexp3}), as well as observable levels of statistical anisotropy \\cite{dklr,abramo,bartolo3} and anisotropic non-gaussianity \\cite{dimopoulos,mythesis,bartolo3} which can serve as discriminators among different inflationary models \\cite{wmap7,groeneboom}.  Indeed, consistency relations among the different levels of non-gaussianity and the level of statistical anisotropy may be obtained for different classes of models so that they may easily be ruled out by observation \\cite{vrl,vr,mythesis,beltran}.  Two of the most notorious differences with respect to the multi-scalar field case of Ref. \\cite{Byrnes:2007tm} are the possibility of parity violating interactions in the action and the possibility of anisotropic expansion; in the former case, the direction of the momentum flow in the Feynman-like diagrams is relevant, something which is not present for scalar and vector boson propagators in the Feynman rules of the particle physics context;  in the latter, the particle production process is statistically anisotropic, which render  the scalar-scalar field perturbation spectra and the vector-vector field perturbation spectra for the different polarizations wavevector-dependent, as well as making the correlators between scalar and vector field perturbations different to zero (producing scalar-vector field perturbation spectra).  In view of recent relevant works where the anisotropic expansion is quite prolonged \\cite{himmetoglu4,watanabe2,dulaney,peloso,watanabe1}, we are urged to comprehensively study the generation of statistical anisotropy in that kind of models. %"
    },
    "1107/1107.1656_arXiv.txt": {
        "abstract": "We investigate the signature of primordial non-Gaussianities in the weak lensing bispectrum, in particular the signals generated by local, orthogonal and equilateral non-Gaussianities. The questions we address include the signal-to-noise ratio generated in the Euclid weak lensing survey (we find the $1\\sigma$-errors for $f_\\mathrm{NL}$ are 200, 575 and 1628 for local, orthogonal and equilateral non-Gaussianities, respectively), misestimations of $f_\\mathrm{NL}$ if one chooses the wrong non-Gaussianity model (misestimations by up to a factor of $\\pm3$ in $f_\\mathrm{NL}$ are possible, depending on the choice of the model), the probability of noticing such a mistake (improbably large values for the $\\chi^2$-functional occur from $f_\\mathrm{NL}\\sim200$ on), degeneracies of the primordial bispectrum with other cosmological parameters (only the matter density $\\Omega_m$ plays a significant role), and the subtraction of the much larger, structure-formation generated bispectrum. If a prior on a standard $w$CDM-parameter set is available from Euclid and Planck, the structure formation bispectrum can be predicted accurately enough for subtraction, and any residual structure formation bispectrum would influence the estimation of $f_\\mathrm{NL}$ to a minor degree. Configuration-space integrations which appear in the evaluation of $\\chi^2$-functionals and related quantities can be carried out very efficiently with Monte-Carlo techniques, which reduce the complexity by a factor of $\\mathcal{O}(10^4)$ while delivering sub-percent accuracies. Weak lensing probes smaller scales than the CMB and hence provide an additional constraint on non-Gaussianities, even though they are not as sensitive to primordial non-Gaussianities as the CMB. ",
        "introduction": "As cosmological data improves, it is becoming increasingly feasible to probe models of the early universe. In particular, primordial non-Gaussianity has emerged as a leading window onto the physics of inflation and the early universe \\citep[for reviews, see][]{2004PhR...402..103B, Komatsu:2010hc, 2011arXiv1102.5052L}. Although non-Gaussian perturbations could in principle take any form, in practice searching for just a few shapes of the bispectrum (3-point correlation function) allows one to discriminate between entire classes of models \\citep{2003PhRvD..67l1301B, Komatsu:2009kd, 2011ApJS..192...18K}. Many alternative models to inflation produce the same non-Gaussian shapes, and therefore can also be constrained at the same time.  The theory of primordial non-Gaussianities is very evolved, and predictions for non-Gaussianities in different observational channels have been made: number counts and large-scale structure statistics \\citep{2010AdAst2010E..64V, 2010CQGra..27l4011D, 2010AdAst2010E..89D, 2011MNRAS.414.1545F}, the cosmic microwave background even outside the Sachs-Wolfe regime \\citep{2007PhRvD..76h3523F, 2010PhRvD..82b3502F} and gravitational lensing \\citep{2011arXiv1103.5396F, 2011MNRAS.411..595P, 2011ApJ...728L..13M, 2011arXiv1104.0926J}.  Most constraints on non-Gaussianities are reported using CMB-observations, either by measuring the bispectrum of the temperature perturbation directly \\citep{2003NewAR..47..797K, 2006JCAP...05..004C, 2008PhRvL.100r1301Y, 2009ApJ...706..399C, 2009ApJS..180..330K, 2009MNRAS.397..837V, 2010MNRAS.404..895V, 2011MNRAS.411.2019C}, by measuring the skewness of weighted averaged CMB-patches \\citep{2004ApJ...613...51M} or by quantifying the corresponding Minkowski functionals \\citep{2005MNRAS.358..684C, 2007MNRAS.377.1668G, 2008MNRAS.389.1439H}. The tightest bounds on the amplitude $f_\\mathrm{NL}$ of the bispectrum, $-10\\leq f_\\mathrm{NL}\\leq74$ has been obtained by \\citet{2011ApJS..192...18K}.  So far, observational constraints on the weak lensing bispectrum \\citep{2003A&A...397..405B, 2011arXiv1105.2309S} mainly concerned the non-Gaussianities generated by structure formation and their breaking of the $\\Omega_m$-$\\sigma_8$ degeneracy \\citep{2003A&A...409..411M}, and first results have been obtained using the skewness of the aperture-mass statistic \\citep{1998MNRAS.296..873S, 2011MNRAS.410..143S}.  In this paper we forecast the constraints which the weak lensing bispectrum will be able to make on primordial non-Gaussianity, especially with a view to Euclid. Although the constraints are not competitive with the CMB if the non-Gaussianities are scale independent, they are complementary since they probe smaller scales compared to the CMB and provide constraints on a possible scale-dependence \\citep{2008JCAP...04..014L}. A scale dependence of $f_\\mathrm{NL}$ at the same order as the spectral index of the power spectrum is natural \\citep{2010JCAP...10..004B} and it may be much stronger, e.g.~a ``step--function\" which is zero on large scales and large on small scales \\citep{Riotto:2010nh}. Weak lensing also has lower systematic errors than other large scale structure probes such as the galaxy bispectrum, scale dependent bias and cluster counts. Because the weak shear provides a linear mapping of the cosmic matter distribution, the statistical properties of the source field are conserved in the observable.  After a brief recapitulation of cosmology and structure formation in Sect.~\\ref{sect_cosmology} we introduce primordial  and structure formation non-Gaussianities in Sect.~\\ref{sect_nong}, in particular the bispectral shapes and the motivation for studying them. The mapping of non-Gaussianities by weak gravitational lensing is treated in Sect.~\\ref{sect_weak_lensing}, where we investigate the properties of the weak lensing bispectrum in its scale- and configuration dependence, and how it builds up as a function of survey depth. Statistical questions concerning the signal strength and misestimations of the non-Gaussianity parameter are addressed in Sect.~\\ref{sect_statistics}, before we focus on systematic errors in the non-Gaussianity parameter due to incompletely removed structure formation non-Gaussianities in Sect.~\\ref{sect_systematics}. We summarise our main results in Sect.~\\ref{sect_summary} and provide visualisations of the weak lensing bispectrum sourced by different non-Gaussianity shapes in Appendix~\\ref{sect_configuration}.  The reference cosmological model used is a spatially flat $w$CDM cosmology with Gaussian adiabatic initial perturbations for the cold dark matter density. The specific parameter choices are $\\Omega_m = 0.25$, $n_s = 1$, $\\sigma_8 = 0.8$, $\\Omega_b=0.04$ and $H_0=100\\: h\\:\\mathrm{km}/\\mathrm{s}/\\mathrm{Mpc}$, with $h=0.72$. The dark energy equation of state is constant in time with a value of $w=-0.9$. ",
        "conclusions": "\\label{sect_summary} The topic of this paper are measurements of primordial bispectra in weak shear data from Euclid, comparisons between different types of non-Gaussianity configuration dependences, statistical questions concerning the inference of the non-Gaussianity parameter $f_\\mathrm{NL}$ and the removal of the much stronger structure formation induced bispectrum. Although not as sensitive as observations of the CMB-bispectrum or the galaxy bispectrum for scale-free non-Gaussianities, weak lensing can place useful independent constraints on non-Gaussianities, in particular on smaller scales where CMB bounds might not apply. It is less prone to systematics than other large-scale structure probes and provides a direct linear mapping of the density field, which conserves its statistical properties.  \\begin{enumerate} \\item{Primordial non-Gaussianities provide a rather weak signal in the weak shear bispectrum \\citep[because of the Gaussianising effect of the line-of-sight integration, see][]{2011arXiv1104.0926J}, and signal-to-noise ratios of order unity can only be expected for $f_\\mathrm{NL} = 200, 575, 1628$ for local, orthogonal and equilateral non-Gaussianities, respectively, where this measurement is most sensitive to scales $0.1\\ldots1~(\\mathrm{Mpc}/h)^{-1}$. These bounds are weaker than those from e.g. observations of the CMB bispectrum, but will serve nevertheless for cross validation, in particular given the absence of strong systematics is weak shear data, or as bounds on scale-dependent non-Gaussianity, because weak lensing maps out scales which are not constrained by the CMB and is sensitive to scales probed by number counts.} \\item{Configuration space integrations can be very efficiently carried out by Monte-Carlo integration schemes, at a fraction of the computational cost. Computations of the signal-to-noise ratio, of $\\chi^2$-functionals or of the Fisher-matrix $F_{\\mu\\nu}$ can be done with accuracies below a percent with $\\mathcal{O}(10^5)$ evaluations instead of $\\mathcal{O}(10^9)$ evaluations for the direct sum over $\\ell_1$, $\\ell_2$ and $\\ell_3$. Very good results were obtained with the CUBA library \\citep{2005CoPhC.168...78H}.} \\item{Fitting the wrong bispectrum type to data yields serious misestimates in the non-Gaussianity parameter $f_\\mathrm{NL}$. Depending on the combination of true and false model there are two cases: either the estimated $f_\\mathrm{NL}$ becomes very small, or the estimate for $f_\\mathrm{NL}$ is a factor of $\\sim\\pm 3$ too large. When looking at numerical values for the $\\chi^2$-functional, one would notice strong discrepancies between data and model when fitting the wrong non-Gaussianity type from values of $f_\\mathrm{NL}$ of a few hundred on.} \\item{The much stronger structure formation bispectrum can be subtracted with a prediction of its bispectrum from perturbation theory if the cosmology is known precisely enough. Propagating the uncertainty in the cosmological parameter set onto the misestimation of $f_\\mathrm{NL}$ if the structure formation bispectrum is not correctly subtracted yielded typical uncertainties of 29, 98 and 149 for local, orthogonal and equilateral non-Gaussianities, much less than the statistical accuracy. As a prior on the cosmological parameters we assumed a Gaussian likelihood for a $w$CDM model combining Euclid's weak shear with baryon acoustic oscillations and Planck's observations of primary CMB anisotropies. Similar ratios between the numerical value of $f_\\mathrm{NL}$ and the standard cosmological parameters were found by \\citet{2011MNRAS.411..595P}.} \\end{enumerate} Many of our investigations can be straightforwardly generalised to other probes of large-scale structure statistics. We intend to generalise our investigations to higher polyspectra and to apply ideas from Bayesian model selection \\citep{2007MNRAS.378...72T,2008ConPh..49...71T} for assigning probabilities to the problem of choosing the correct non-Gaussianity type."
    },
    "1107/1107.3838_arXiv.txt": {
        "abstract": "We analyze \\galex\\ UV data for a system of four gravitationally-bound groups at $z=0.37$, SG1120, which is destined to merge into a Coma-mass cluster by $z=0$, to study how galaxy properties may change during cluster assembly. Of the 38 visually-classified S0 galaxies, with masses ranging from \\hbox{$\\log(M_*)[\\msun]\\approx10$--$11$}, we detect only one in the NUV channel, a strongly star-forming S0 that is the brightest UV source with a measured redshift placing it in SG1120. Stacking the undetected S0 galaxies (which generally lie on or near the optical red-sequence of SG1120) still results in no NUV/FUV detection ($<2\\sigma$). Using our limit in the NUV band, we conclude that for a rapidly truncating star formation rate, star formation ceased {\\it at least} $\\sim0.1$ to $0.7$~Gyr ago, depending on the strength of the starburst prior to truncation. With an exponentially declining star-formation history over a range of time-scales, we rule out recent star-formation over a wide range of ages. We conclude that if S0 formation involves significant star formation, it occurred well before the groups were in this current pre-assembly phase. As such, it seems that S0 formation is even more likely to be predominantly occurring outside of the cluster environment. ",
        "introduction": "\\label{sec:intro}  S0 galaxies are more common in denser environments than in the field \\citep{Dressler80}, and the fraction of S0 galaxies increases over time \\citep{Dressler97,Fasano00,Desai07}, such that groups/clusters at $z\\sim0$ have S0 fractions $\\approx3$ times greater than at $z\\sim0.5$. Due to the commensurate decline in the spiral fraction, these findings have been interpreted as arising from the transformation of spirals into S0's. Further observations have refined the model to suggest that over this redshift range groups are the primary site of S0 formation \\citep[e.g.,][]{Wilman09,Just10}, i.e. the galaxies are ``preprocessed'' in groups prior to accretion into the cluster \\citep[e.g.,][]{Zabludoff96}.  However, the correlation between groups and clusters, and the uncertainty in determining whether one is observing a group that will soon fall into a cluster, complicate the interpretation of environmentally dependent evolution. After all, galaxy properties begin to change well outside of what is typically referred to as a cluster \\citep[i.e., 2 to 3 virial radii;][]{Lewis02,Gomez03}. The question then becomes whether S0 formation occurs in isolated groups or only when a group enters this meta-cluster environment. Is S0 formation related to the changes in star formation properties observed in the far outskirts of clusters?  Super Group 1120-1202 (hereafter SG1120) provides a unique opportunity to address these questions. It is a bound collection of four galaxy groups at $z\\sim0.37$ that is in the process of assembling into a cluster. The four groups will merge by $z=0$ to form a cluster one-third the mass of Coma or greater \\citep{Gonzalez05}, yet they are clearly independent groups as observed. Spectroscopic redshifts and morphological classifications exist, allowing detailed analysis of its constituent galaxies. The fraction of S0 galaxies in SG1120 is already as high as that of clusters at similar redshift \\citep{Kautsch08}, demonstrating that the high-density, massive cluster environment is not the primary site of S0 formation. The question of whether these S0's formed recently, in the pre-assembly epoch, is that which we now consider.  To determine whether the S0's formed recently, we measure their recent star formation history (SFH). A host of different mechanisms have been suggested for the transformation, including mergers and galaxy-galaxy interactions \\citep{Toomre72,Icke85,Lavery88,Byrd90,Mihos04,Bekki11}, ram pressure stripping \\citep{Gunn72,Abadi99,Quilis00}, strangulation \\citep{Larson80,Bekki02}, and harassment \\citep{Richstone76,Moore98}. The different mechanisms have their own strengths and weaknesses. A difficulty with ram-pressure stripping as the primary mechanism lies with accounting for the large fraction of S0's in the field \\citep[e.g.,][]{Dressler80,Dressler04}, although ram-pressure stripping has been clearly observed in clusters \\citep[e.g.,][]{Irwin87,Kenney99} and could account for the deficiency of HI gas observed in cluster spirals \\citep[e.g.,][]{vandenBergh76,Giovanelli83,vanGorkom96,vanGorkom04}. On the other hand dynamical interactions (i.e. mergers and tidal effects) are consistent with groups as the primary site of S0 formation \\citep[e.g.,][]{Wilman09,Just10}, although in this scenario it is unclear why a comparable star-formation quenching efficiency is observed in both groups and clusters \\citep{Poggianti09b}. For an excellent review of these different mechanisms and their ability to explain observations across different environments and redshift, we refer the reader to \\cite{Boselli06}. These processes all involve the halting of star formation, but operate on different timescales and affect the SFH differently. By focusing on the SFH's of the S0's in SG1120 we can constrain these mechanisms acting in a currently assembling cluster.  Some measures of the SFH's of the S0's in SG1120 are already available. Nearly all the S0's lie on or near the optical $B-V$ red sequence (Figure~1) and inspection of their optical spectra reveal no emission lines, suggesting no significant ongoing star formation. Strong Balmer absorption indicative of star-formation within the past $\\sim1$~Gyr (so-called E$+$A galaxies; initial work by \\cite{Dressler83} and recent work, e.g., \\cite{Yang08}) is also absent in their spectra. However, all of these signatures are primarily sensitive to significant bursts of recent star formation ($\\sim$ tens of percents by stellar mass). If the S0 formation process involves more modest bursts (or just a truncation of a low level of star formation), and if this happened recently ($<1$~Gyr ago), then detection in the UV may be the best way to identify it. With these goals in mind, we have obtained GALEX \\citep{Martin05,Morrissey05} imaging of SG1120.  Our paper is organized as follows. In \\S2 we describe the data that appear in this study. In \\S3 we present our results, which we then discuss and summarize in \\S4. We adopt a cosmology with $H_{0}=70$~\\kms~Mpc$^{-1}$, $\\Omega_{m}=0.3$, and $\\Omega_{\\Lambda}=0.7$. Optical magnitudes are in the Vega system while UV magnitudes are in the AB system; one can convert the $B$ and $V$ magnitudes to the AB system by adding $-0.275$ and $-0.116$, respectively.  \\begin{figure} \\epsscale{1.0} \\plotone{figure1.eps} \\figurenum{1} \\caption{\\label{optCMD} Rest-frame $B-V$ color magnitude diagram for spectroscopically confirmed SG1120 galaxies. S0 galaxies are highlighted as stars and the remaining members are shown as circles. NUV-detected galaxies are marked using filled-symbols, and approximate tracks of constant stellar mass are overplotted (see \\S~\\ref{sec:data}). Most S0's lie on the red sequence, consistent with being dominated by an old, passively evolving stellar population, and comprise $\\sim35\\%$ of all red sequence galaxies.} \\end{figure} ",
        "conclusions": "\\label{sec:discuss}  Our chief finding is a lack of NUV emission in the S0 galaxies in SG1120, down to $m_{NUV}=26.0$, or $0.01$~\\msun~yr$^{-1}$. Evidently the S0's with masses from $\\log(M_*)[\\msun]\\approx10$--$11$ are not forming many new stars, but the time since their last significant star forming episode depends on their SFH. Generally, if star formation shut off rapidly, then they could have formed stars more recently. Conversely, if their star formation turned off gradually, or if they experienced a significant burst of star-formation prior to the shut-off, then more time must have passed for them to drop below our detection threshold. We investigate both possibilities.  \\begin{figure*} \\epsscale{0.8} \\plotone{figure5.eps} \\figurenum{5} \\caption{\\label{extsfh} Plot of e-folding time ($\\tau$) vs.  $t_{\\rm thresh}$ for model galaxies with exponentially declining SFH's, where $t_{\\rm thresh}$ is the time for the model galaxy to drop below our NUV detection threshold.  The model galaxies have $\\log(M_*)[\\msun]=10.5$ at the time their SFR begins to decline. The shaded regions, from darkest to lightest, are ruled out by our NUV detection limit assuming SFR's of $1$, $5$, and $10$~\\msun~yr$^{-1}$, respectively. The dashed lines demarcate the parameter space considered in \\cite{Balogh11}. } \\end{figure*}  In the rapid truncation scenario (Figure~\\ref{scenarios}), our models show that the minimum time since the burst ranges from $\\sim0.1$ to $0.7$~Gyr, depending on the mass of the galaxy and the strength of the burst. While these models are consistent with the S0's having formed at much earlier times ($>1$~Gyr), in the ``no-burst'' model the S0's could have stopped forming stars as recently as $0.1$--$0.2$~Gyr ago, depending on the mass. In other words, if the formation of an S0 involves a morphological transformation and a halting of star formation, but no additional star formation, we cannot use UV photometry to constrain meaningfully the time since that event. However, if their formation involved an episode of significant star formation, as one might expect in a merger, then they must have stopped forming stars $>0.3$~Gyr prior to the time at which we are observing them. Given the lack of E+A spectra among our S0's, which indicate star-formation within the past $\\sim1$~Gyr, it is likely that the S0's formed at even earlier times. While the strength of the absorption will be weaker for galaxies with no burst of star-formation, \\cite{Yang08} find E+A galaxies with burst fractions as low as $7\\%$ are consistent with their observations, demonstrating that low burst strengths can still yield measurable E+A spectra.  We next consider the limits we can place on a gradual reduction in SFR.  \\cite{Moran07} find evidence for newly-formed S0's in groups at the outskirts of two massive clusters at $z\\sim0.5$. In the process of forming, the SFR's of these S0's is interpreted to consist of a gradual decline over a $\\sim1$~Gyr timescale, consistent with strangulation. In Figure~\\ref{extsfh}, we model the S0's in SG1120 with similarly extended SFH's. Our S0's are consistent with a similar slow conversion, provided that they started this decline at earlier times. The current phase in SG1120's evolution, as the four groups merge, is therefore unlikely to play the dominant role in S0 formation. Interestingly, a population of galaxies that lie in the so-called ``green valley'' have been identified in groups at $z\\sim0.8$--$1$ \\citep{Balogh11}, and have been interpreted as those moving from the blue cloud to the red sequence due to an exponentially declining SFR with $\\tau\\sim0.6$--$2$. These galaxies are candidate S0 progenitors given (1) their presence in groups, and (2) their intermediate colors, since (red) S0's forming from (blue) spirals must traverse a similar path in color space. Our models in Figure~\\ref{extsfh} have stellar masses typical of these transition candidates. Models with an initial SFR of $1$~\\msun~yr$^{-1}$ over the full range of $\\tau\\sim0.6$--$2$ are consistent both with these high redshift objects and our UV limits. Models with higher initial SFR's begin to violate our limits for certain combinations of $\\tau$ and $t_{\\rm thresh}$. Galaxies similar to these ``green valley'' group galaxies could be the progenitors of the S0's in SG1120, but this would again imply that the cluster assembly process is not associated with the S0 transformation phase.  A similar picture appears to unfold at $z=0$. \\cite{Hughes09} find locally that ``green valley'' galaxies in $NUV-H$ color are predominantly HI-deficient spirals with quenched star formation found in higher-density environments. Further analysis has shown that these galaxies are consistent with migration from the blue cloud to the red sequence over at least a $\\sim3$~Gyr timescale due to ram-pressure stripping \\citep{Cortese09}. A concordant result is also found over a wider range of density \\citep{Gavazzi10}.  While these results at low redshift cannot be directly applied to higher $z$, they demonstrate that a slow process of migration across the ``green valley'' is a viable physical mechanism for quenching star formation, which for SG1120 would require S0 formation prior to the cluster assembly phase.  Although the S0 fraction of SG1120 is already sufficiently large to match that of Coma within the uncertainties and the scatter in S0 fractions, one could envision the S0 fraction of SG1120 growing by as much as a factor of two between its current redshift and today. If so, then S0's should be added at a rate of $\\sim3$--10 per Gyr. For models with a gradual halting in the SFR, this implies that a significant number of S0's should be in the process of forming. However, the likely progenitors candidates are not seen: there are $\\sim 6$ non-star-forming ``passive spirals'', and at most one star-forming S0. Hence, if strangulation is chiefly responsible for S0 formation, then SG1120 has finished forming S0's. Conversely, if S0 formation is ongoing in this system, then the S0's are forming without a gradual reduction in SFR \\citep[e.g.,][]{vandenBergh09}.  We find that nearly all of the S0's in SG1120 show no trace of star-formation, and by modeling their star-formation histories with both a rapid truncation and a gradual reduction in SFR, are able to place limits on the time since their last significant star-forming episode. Our constraints are weaker in the rapid reduction scenario, particularly if S0 formation does not involve a significant burst of star formation; from our models, the S0's could have formed stars as recently as $\\sim0.1$~Gyr ago and be consistent with our NUV limit. In models where a burst of star-formation occurs, forming at least $20\\%$ of the stellar mass, our limits imply that this occurred at least $\\sim0.3$~Gyr ago. If a more gradual reduction in star-formation occurred, modelled as an exponentially declining SFR from a level of $1$--$10$~\\msun~yr$^{-1}$, then our limits increase to $\\sim$ several Gyr. This scenario is incompatible with SG1120 continuing to form new S0's, as a significant number of transition galaxies would be expected that are not observed. Evidently, the formation of S0's occurred prior to the assembly phase of the cluster."
    },
    "1107/1107.4629_arXiv.txt": {
        "abstract": "Current and upcoming wide-field, ground-based, broad-band imaging surveys promise to address a wide range of outstanding problems in galaxy formation and cosmology.  Several such uses of ground-based data, especially weak gravitational lensing, require highly precise measurements of galaxy image statistics with careful correction for the effects of the point-spread function (PSF).  In this paper, we introduce the {\\sc shera} (SHEar Reconvolution Analysis) software to simulate ground-based imaging data with realistic galaxy morphologies and observing conditions, starting from space-based data (from COSMOS, the Cosmological Evolution Survey) and accounting for the effects of the space-based PSF.  This code simulates ground-based data, optionally with a weak lensing shear applied, in a model-independent way using a general Fourier space formalism. The utility of this pipeline is that it allows for a precise, realistic assessment of systematic errors due to the method of data processing, for example in extracting weak lensing galaxy shape measurements or galaxy radial profiles, given user-supplied observational conditions and real galaxy morphologies.  Moreover, the simulations allow for the empirical test of error estimates and determination of parameter degeneracies, via generation of many noise maps. The public release of this software, along with a large sample of cleaned COSMOS galaxy images (corrected for charge transfer inefficiency), should enable upcoming ground-based imaging surveys to achieve their potential in the areas of precision weak lensing analysis, galaxy profile measurement, and other applications involving detailed image analysis. ",
        "introduction": "\\label{S:intro}  A tremendous variety of measurements are carried out on astronomical images from ground-based telescopes.  A generic problem that often arises is the question of how the intrinsically limited resolution of ground-based images (both due to convolution with the atmospheric point-spread function, or PSF, and due to the finite pixel size) affects our ability to measure quantities such as the radial profiles of galaxies $I(r)$, or statistics of the profile such as its slope, half-light radius, and ellipticity.  Moreover the error distributions of these parameters, which are often estimated via a highly non-linear fitting procedure, are also unclear in many circumstances.  One application that particularly suffers from such uncertainties is the field of weak gravitational lensing (for a review, see \\citealt{2001PhR...340..291B}, \\citealt{2003ARA&A..41..645R}, \\citealt{2008ARNPS..58...99H}, or \\citealt{2010RPPh...73h6901M}).  In the past decade, weak lensing has been used extensively for measurements that can constrain cosmology and galaxy formation. Cosmic shear measurements have constrained cosmological parameters (e.g., most recently, \\citealt{2008A&A...479....9F,2010A&A...516A..63S}); cluster lensing analyses (e.g., \\citealt{2007MNRAS.379..317H,2010PASJ...62..811O}) have been used to understand the most massive structures in the universe, the abundance of which can constrain cosmology through the mass function \\citep[e.g., ][]{2007ApJ...657..183R,2008MNRAS.387.1179M,2009ApJ...692.1060V,2010ApJ...708..645R}; and galaxy-galaxy lensing measurements have probed the connection between galaxies, their dark matter halos, and their larger scale environments \\citep{2010MNRAS.408.1463S,2011arXiv1104.0928L}, as well as constraining the theory of gravity on cosmological scales when combined with other observational methods \\citep{2010Natur.464..256R}. The next decade promises a larger volume of weak lensing data that can be used for more precise constraints on cosmology and galaxy formation, from surveys such as Hyper Suprime-Cam (HSC, \\citealt{2006SPIE.6269E...9M}), Dark Energy Survey (DES\\footnote{\\texttt{https://www.darkenergysurvey.org/}}, \\citealt{2005astro.ph.10346T}), the KIlo-Degree Survey (KIDS\\footnote{\\texttt{http://www.astro-wise.org/projects/KIDS/}}), the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS\\footnote{\\texttt{http://pan-starrs.ifa.hawaii.edu/public/}}, \\citealt{2010SPIE.7733E..12K}); and even more ambitious programmes are planned such as the Large Synoptic Survey Telescope (LSST\\footnote{\\texttt{http://www.lsst.org/lsst}}, \\citealt{2009arXiv0912.0201L}), Euclid\\footnote{\\texttt{http://sci.esa.int/science-e/www/area/index.cfm?fareaid\\ =102}}, and the Wide-Field Infrared Survey Telescope (WFIRST\\footnote{\\texttt{http://wfirst.gsfc.nasa.gov/}}).  Weak lensing measurements depend on precise measurements of the shapes of galaxies, in an attempt to infer the apparent shape distortions in distant `source' galaxies due to the mass in nearby `lenses' (galaxies, clusters, or other mass distributions). Weak lensing is a statistical measurement, with averages over large numbers of sources in order to detect the $\\sim 0.1$--$1$ per cent level distortions within the noise of the intrinsic galaxy ellipticities, which are typically a factor of $\\sim 100$ larger.  Unfortunately, coherent systematic distortions of galaxy shapes due to the PSF caused by the atmosphere (for a ground-based telescope), telescope optics, and detector are significantly larger than typical weak lensing distortions, which means that accurate PSF-correction is critical for current lensing studies, and all the more so for future lensing surveys that aim for $<1$ per cent statistical errors.  The weak lensing community has had several challenges, using blind simulations, to identify the most promising methods of PSF correction. The first of these, the Shear TEsting Programme-1 \\citep[STEP1;][]{2006MNRAS.368.1323H}, included mock galaxies with idealized radial profiles and several PSFs meant to mimic specific observational issues (e.g., astigmatism).  The second, STEP2 \\citep{2007MNRAS.376...13M}, used shapelets \\citep{2003MNRAS.338...35R,2003MNRAS.338...48R} decompositions of COSMOS galaxies (including the COSMOS PSF, for which no correction was imposed) as inputs, and then convolved them with a variety of ground-based PSFs from Subaru Suprime-Cam \\citep{2002PASJ...54..833M}.  Finally, the GRavitational lEnsing Accuracy Testing-08 \\citep[GREAT08;][]{2009AnApS...3....6B,2010MNRAS.405.2044B} and GREAT10 \\citep{2010arXiv1009.0779K} challenges reverted to composite model galaxies (based on S\\'ersic profiles; \\citealt{1968adga.book.....S}) with very specific sets of parameters, and tested the recovery of the shear as a function of image $S/N$, PSF size, and galaxy profile type.  All of these challenges were useful to the lensing community in different ways, and in some cases led to changes in attitudes towards (or a greater understanding of) common methods of PSF correction.  However, their ability to identify, in a broad sense, the most promising methods of PSF correction does not mean that they can be used to constrain, to high precision, the shear calibration in \\newtext{all lensing observations} using those PSF-correction methods\\footnote{\\newtext{The validity of this statement depends on the data for which the shear calibration is desired.  For example, the STEP2 simulations are likely to give a more realistic estimate of the calibration for the Subaru Suprime-Cam data that it was designed to mimic than for data from some other telescope (assuming that image combination and other steps in a realistic data analysis, which were not simulated, do not introduce additional biases).}}.  There are numerous reasons why this is the case, such as galaxy models and observing conditions (PSF and depth) that are not representative of a particular survey. STEP2 \\citep{2007MNRAS.376...13M} showed that the shear systematics for nearly every method of PSF correction depend on the observing conditions (the galaxy $S/N$ and resolution compared to the PSF); GREAT08 \\citep{2010MNRAS.405.2044B} showed a dependence on galaxy type as well.  \\newtext{More recent work has shown that the dependence of shear calibration factors on the galaxy population is generic. In particular, \\citet{2007MNRAS.380..229M} and \\citet{2011MNRAS.414.1047Z} showed that there is {\\em no} stable shape measurement algorithm on finite-resolution data whose shear calibration factor is independent of the galaxy population\\footnote{The argument hinges on the existence of a finite number of well-measured moments $M_{ij}$ of the galaxy, and the fact that the dependence of the $M_{ij}$ on shear ($\\partial M_{ij}/\\partial\\gamma_k$) is determined in part by higher, unmeasured moments. \\reftext{A related issue occurs in Fourier space: when attempting to define a roundness test for a sheared galaxy, Bernstein (2010, last paragraph of section 4.2) finds that the finite extent of the modulation transfer function (MTF) prevents such a test from being shear-covariant, and argues that the issue is generic.}}. Instead, \\citet{2010MNRAS.406.2793B} and \\citet{2011MNRAS.414.1047Z} argue that the same lensing survey that measures shear could also determine the range of galaxy models presented to us by the real Universe. \\reftext{While we will explore this point in more detail in Sec.~\\ref{S:psfcorrection},} we conclude that to precisely constrain the shear calibration or understand observational selection biases in any given survey, the simulations must have realistic galaxy models as well as observing conditions.}  When constraining systematic errors in ground-based lensing data, we may wish to use space-based data as the basis of our simulations, due to its much higher resolution.  Indeed, one might ask why simulations are needed at all: can we simply rely on comparison with PSF-corrected shapes on space-based images?  This approach was taken by \\cite{2008ApJ...684...34K}, who compared shape measurements using a KSB-based method \\citep{1995ApJ...449..460K,2003ApJ...597...98H} on Subaru data against shape measurements using the RRG method \\citep{2000ApJ...536...79R} on COSMOS data. However, there are some limitations to this approach.  First, \\newtext{if one wishes to avoid complications due to a non-negligible PSF and therefore the need for substantial PSF correction in the space-based data, the comparison must be restricted to a subset of larger galaxies, as in \\cite{2008ApJ...684...34K}.} \\newtext{More importantly}, such a direct comparison of the derived shapes or shears is not possible at all for some shape measurement methods if the ellipticities are defined in incompatible ways, as will be explained at greater detail later in this paper \\newtext{(Sec.~\\ref{SS:shapecatalog})}.   We thus conclude that, rather than carrying out a catalogue-level comparison, we should use the space-based data to make realistic simulations of ground-based lensing data, to which a shear can be added and shear recovery can be tested.  In this paper, we therefore have three goals.  First, we outline a method for simulating ground-based lensing data using higher-resolution data from space, including a careful treatment of the original space-based PSF and pixel sampling, and conversion to the new ground-based PSF and pixel sampling, inspired by \\cite{2000ApJ...537..555K}.  This method will allow the galaxies to be sheared, so that we can test the recovery of gravitational shear, including the many types of selection biases and PSF effects such as those described in \\cite{2004MNRAS.353..529H} and \\cite{2005MNRAS.361.1287M}.  This method is a model-free generalization of that used in \\cite{2010PASP..122..947D}, which relies on shapelets decompositions (Gauss-Laguerre basis functions) and therefore depends on the galaxy profiles and PSFs being well-described by sums of these functions to some finite order. Given that  this assumption is not necessarily true for realistic galaxies and PSFs (e.g., \\citealt{2010A&A...510A..75M}), the advantage of a model-independent image simulation method is clear. Second, we describe a publicly-released implementation of this method in IDL.  We emphasize that, while this paper focuses on weak gravitational lensing, this simulation pipeline is equally applicable to many other science analyses that are commonly done using ground-based data, for example modeling of galaxy radial profiles.  Finally, we demonstrate the method by simulating Sloan Digital Sky Survey (SDSS) lensing data, and show how these simulations can be used to estimate shear systematics in SDSS to high precision.  This will be of practical use for interpreting past lensing measurements that used the SDSS shape catalogue we simulate, and eventually for reducing the systematic error budgets in future papers using that catalogue.  We begin in Section~\\ref{S:lensing} by describing weak gravitational lensing and its effects on galaxy shapes. Section~\\ref{S:psfcorrection} describes the process of removing the effects of the PSF from measured galaxy shapes, so that lensing can be measured.  The space-based data that are used as the basis of these simulations, and the ground-based data that we simulate in this paper, are described in Section~\\ref{S:data}.  In Section~\\ref{S:methodology}, we describe the methodology that will be used to create accurate simulations of ground-based data.   We describe steps that we took to prepare the space-based data for this purpose in Section~\\ref{S:imageprep}, and our specific implementation of the simulation method in Section~\\ref{S:implementation}. \\reftext{Technical tests of this implementation are presented in Section~\\ref{S:techvalidation}, and an example of how the {\\sc shera} outputs can be used to test galaxy shape measurements is in Section~\\ref{S:results}. } We discuss these results in Section~\\ref{S:discussion}. ",
        "conclusions": "\\label{S:discussion}  We have described new, publicly-available software, {\\sc shera}, that can be used to simulate (optionally sheared) ground-based images with realistic morphologies and any PSF that has worse resolution compared to COSMOS. This software is independent of modeling assumptions about what galaxies or PSFs look like, properly handles the pixel response functions, and has been tested to sub-per cent precision \\reftext{(Sec.~\\ref{S:techvalidation})}.  The code has been publicly released, along with CTI-corrected, cleaned COSMOS galaxy postage stamps for a flux-limited sample at $F814W<22.5$.  This code should be highly useful for realistically assessing the systematics of lensing analysis (or indeed any other detailed image analysis, such as galaxy profile determination) in ground-based data, including the effects of a range of observing conditions.  It should also allow for assessment of parameter uncertainties and degeneracies, and selection biases, via the generation of many noise realizations for a single galaxy.  As a basic demonstration, we have shown a crude estimation of shear calibration bias for SDSS single-epoch lensing analysis.  In order for the code and data that are described in this paper to be more broadly applicable, there are several improvements that would be needed.  First, we would like the ability to simulate data in other passbands besides $i$ \\newtext{(for which there is precedent, e.g. \\citealt{2008APh....30...65F} used data from the Ultra Deep Field to do so)}.  Realistic galaxies have colour gradients that will cause an intrinsically different appearance in significantly different bands.  While this may be unimportant for current surveys that have statistical errors that are $>5$ per cent, future surveys such as LSST that aim for better than $1$ per cent precision in the lensing signal have correspondingly a need for very well-constrained shear systematics.  Thus, such effects must be handled realistically, and it is possible that sufficient well-sampled data\\footnote{Data with some instruments may not have a Nyquist-sampled PSF, depending on the choice of dither pattern.} in other bands and in random fields exists in the {\\it HST} archive that could be used for this purpose.  We specify random fields because those that were selected due to, e.g., presence of a galaxy cluster, may not have a representative morphological mix.  There is additional value to obtaining data in other, random fields, given that even with the relatively large (for {\\it HST}) size of the COSMOS field, it still exhibits significant cosmic variance (e.g., \\citealt{2010ApJ...708..505K}) in the redshift distribution, which may manifest at some level as an atypical morphology distribution.  Additionally, we need a way to simulate a deeper sample.  Presently we have stopped at $I<22.5$ because the method described here requires modification if the input images have non-negligible noise, and the COSMOS data have $S/N>50$ for that magnitude range.  Thus, we require either a generalization of the method to lower $S/N$ in the input images, or we require deeper input data than is available in the COSMOS field. \\newtext{In some cases, the former would be possible: since the {\\sc shera} algorithm is a linear operation on the COSMOS postage stamps, one could propagate any noise covariance matrix through the pipeline and arrive at an output covariance matrix $N^{\\rm IN}_{ij}$ on the output postage stamps. If the noise covariance in the data we wish to simulate is $N^{\\rm OUT}_{ij}$, then so long as ${\\mathbfss N}^{\\rm OUT}-{\\mathbfss N}^{\\rm IN}$ is semi-positive definite, one could add in the appropriate additional noise and thereby extend the methodology of this paper to noisy input data. The software implementation of such a method, exploration of its range of applicability, and investigation of complications such as non-Gaussian noise is left to a future paper. In the opposite case, namely that where the input data has more noise in some mode than the data we wish to simulate, it seems likely that the problem is hopeless and deeper input data would be required.}  Despite the need for future work to make the code and/or input dataset as useful as possible for lensing surveys that are coming up on the timescale of $\\sim 1$ decade, we anticipate that {\\sc shera v1.0} has numerous applications for better understanding of current data and those surveys that are starting in the next year, such as KIDS, HSC, and DES."
    },
    "1107/1107.0070_arXiv.txt": {
        "abstract": "We have analyzed a \\suzaku long-look of the active galactic nucleus \\mcg, a type 1.5 Seyfert with very little X-ray absorption in the line of sight and prominent features arising from reflection off circumnuclear material: the Fe line and Compton reflection hump. We place tight constraints on the power law photon index ($\\Gamma$=1.80$\\pm$0.02), the Compton reflection strength ($R$=0.69$\\pm$0.05), and the Fe K emission line energy centroid and width ($E$=6.40$\\pm$0.02 keV, $v_\\textrm{FWHM} <$\\,7100 km\\,s$^{-1}$).  We find no significant evidence for emission from strongly ionized Fe, nor for a strong, relativistically broadened Fe line, indicating that perhaps there is no radiatively efficient accretion disk very close in to the central black hole.  In addition we test a new self-consistent physical model from Murphy \\& Yaqoob, the ``\\mytorus'' model, consisting of a donut-shaped torus of material surrounding the central illuminating source and producing both the Compton hump and the Fe K line emission.  From the application of this model we find that the observed spectrum is consistent with a Compton-thick torus of material (column density \\NH=3.6\\err{1.3}{0.8}$\\times 10^{24}$ cm$^{-2}$) lying outside of the line of sight to the nucleus, leaving it bare of X-ray absorption in excess of the Galactic column.  We calculate that this material is sufficient to produce all of the Fe line flux without the need for any flux contribution from additional Compton-thin circumnuclear material. ",
        "introduction": "\\mcg\\ is an X-ray bright Seyfert 1.5 active galactic nucleus (AGN) located at a redshift of $z$ = 0.04686. Past X-ray observations of this source performed with \\textsl{EXOSAT}, \\textsl{ASCA}, \\xmm and \\sax have revealed the following spectral components in addition to the primary X-ray power law: a soft excess, Fe emission lines, and a Compton reflection hump. Importantly, there has been no evidence for X-ray absorption by gas along the line of sight in excess of the Galactic column, indicating that in the X-ray band this AGN is a ``bare nucleus''. This combination makes \\mcg\\ an interesting target of study, since the lack of significant X-ray absorption provides a clean view of the nucleus with a relatively simple spectum to model, while the presence of strong reflection components allows us to place constraints on the physical geometry of the circumnuclear material surrounding the AGN.  Ghosh \\& Soundararajaperumal (1992) analyzed \\textsl{EXOSAT} data obtained in 1984 that revealed the highly variable soft excess below about 2 keV.  They modeled this component with a steep power law in addition to their continuum power law. \\textsl{ASCA} data covering 2.5--10 keV with good CCD resolution were analyzed by Weaver \\etal (1995). They modeled the spectrum using a hard X-ray power law ($\\Gamma = 1.75 \\pm 0.05$) with Galactic absorption and confirmed the need for a soft excess, as well as an Fe \\Ka\\ emission line which was unresolved. Weaver, Gelbord \\& Yaqoob (2001) analyzed two additional \\textsl{ASCA} observations of \\mcg, tracking the Fe line flux over a time scale of years and showing large variation in the flux of the underlying continuum. However large uncertainties precluded definitive conclusions about the variation of the Fe line parameters.  A more recent analysis by Bianchi \\etal (2004) using simultaneous data from \\xmm and \\sax covered a much broader energy range than previous observations (0.5--200 keV).  Unfortunately only 7 ks of good EPIC-pn data were obtained for the source, providing only loose constraints in the Fe K bandpass (the Fe \\Ka\\ line was unresolved with $\\sigma < 340$ eV and an equivalent width, EW, of 45\\err{85}{24} eV). They were able to loosely constrain the Compton reflection hump with $R = 0.4 \\pm 0.3$ and $\\Gamma = 1.72$\\err{0.08}{0.06}.  In this paper we present an in-depth analysis of a single 140 ks long-look \\suzaku observation of \\mcg. \\suzaku is a Japanese observatory that provides broadband X-ray spectra from $\\sim$\\,0.5 keV to above 500 keV with good energy resolution and effective area around 5--7 keV for detailed analysis of the Fe K complex (Mitsuda \\etal 2007). Our goals for this observation were to study the Fe K emission, constrain the Compton hump, confirm the bare nucleus, and study the soft excess, all of which \\suzaku is capable of doing well. This information can then be used to explore possibilities for the geometry of the circumnuclear material including the Fe line emitting gas and we improve upon parameter values in the literature for the Compton hump and Fe K emission lines. This paper is structured as follows: Section 2 details data reduction, Section 3 describes spectral fitting and analysis, and Section 4 contains a discussion of the results.    \\begin{figure} \\plotone{figure1.ps} \\caption{The lightcurve over the duration of the observation.  Bins are 800\\,s.  The XIS count rates are in the 2--10 keV range while the PIN count rates are in the 20--50 keV range.  The increases in flux seen in the XIS around 30--50 ks and 110--140 ks are small (only about 20\\%) and flux-resolved spectroscopy did not reveal significant change in the shape of the spectrum during these episodes.} \\label{figlc} \\end{figure} ",
        "conclusions": "\\subsection{The Fe K Complex} Focusing on the Fe K band we found the need for both Fe \\Ka\\ and \\Kb\\ emission lines as well as an Fe K shell absorption edge (in broadband fits this edge was modeled sufficiently by the edge associated with the Compton reflection hump in both \\pexrav and \\mytorus). From the value of the emission line width found in our best-fit broadband model we calculated the velocity full width at half maximum ($v_\\textrm{FWHM}$) of the emitting material to be $<$\\,7100 km\\,s$^{-1}$.  This is consistent with values obtained for the optical H$\\beta$ broad emission line of around 6400--8500 km\\,s$^{-1}$ (Osterbrock 1977; Kollatschny \\& Dietrich 2006; Winter \\etal 2010) and is a significant improvement on previous upper limits set by Weaver \\etal (1991) and Bianchi \\etal (2004) of $\\lesssim$\\,30,000 km\\,s$^{-1}$. Using a black hole mass estimated from optical luminosity and line widths to have a value of $10^{8.4}$ \\Msun (Bian \\& Zhao 2003; Winter \\etal 2010) and assuming Keplerian motion of the emitting material, we estimated the radius of the emitting region to be\\ $\\gtrsim$\\,45 lt-days or roughly 1200 $R_{\\rm S}$.  We also tested for a broad line and Compton shoulder.  According to de la Calle P\\'erez \\etal (2010), roughly $\\gtrsim$\\,1.5$\\times 10^5$ counts in the 2--10 keV band at CCD resolution provide good enough statistical quality to significantly detect a broad line.  In the combined FI XIS we have $\\sim$\\,4$\\times 10^5$ counts, and our upper limit on the EW of a broad line places us in the lower part of the EW range of detected broad lines in the FERO sample of Seyferts observed with \\textit{XMM-Newton}, wherein significant detections of broad lines with EW's in the range of 50--250 eV were reported. We conclude that a very strong broad line ($\\gtrsim$\\,50 eV) does not exist in MCG--2-58-22 or else our observation would have been sufficient to significantly detect it; if there does exist a broad line in this source, then it must be very weak.  We would expect to see a Compton shoulder given the presence of Compton-thick material, however we did not obtain a significant detection.  It should also be noted that the Compton shoulder is included in the \\mytorus model automatically, based on the strength of the Fe line and the column density of the torus.  \\subsection{Reflection and Geometry of the Circumnuclear Material} This source also shows a very prominent Compton hump around 20--30 keV (see Figure \\ref{figspec}b).  This feature, arising from Compton scattering of high energy photons off Compton thick material in the vicinity of the black hole, is often associated with the Fe line emission seen in AGN, since the same Compton thick material that produces the Compton hump also produces Fe K emission. By knowing the (model-dependent) relationship between the strength of the Compton hump ($R$) and the expected Fe line EW, we can test if Compton thick gas is capable of accounting for the entire observed Fe line flux.  Based on calculations done by George and Fabian (1991) for a disk geometry, we found that the expected Fe line EW for our value of $R$ (assuming an inclination of 30$\\degr$ and solar abundances) is $\\sim$\\,80$\\pm 6$ eV.  In our broadband fits we see a considerably milder Fe line flux with an EW closer to 40 eV, about half of the expected value. We found a similar, though less robust result using \\mytorus, with $A_{\\rm L} =$ 0.75$\\pm 14$ and EW = 47$\\pm 6$ eV (calculated from the flux of the line) which assumes a ``donut-like'' geometry, in this case a Compton-thick torus out of the line of sight to the nucleus. In both cases we have found a lower than expected Fe line EW, possibly due to an underabundance of Fe. This makes it unlikely that there is significant contribution to the Fe line flux from Compton thin material in the vicinity of the black hole. This is in contrast to many Seyferts; for example the sample of Rivers \\etal (2011), which utilized the \\pexrav disk reflection model, found that on average only $\\sim$30\\% of the Fe line flux in Seyferts is associated with the Compton reflection component, implying the presence of substantial amounts of Compton-thin material and/or supersolar abundances of Fe. However other samples of Seyferts observed with \\sax and \\xmm (Perola \\etal 2002; Bianchi \\etal 2004) were shown to have values of EW and $R$ consistent with the Fe line and reflection hump arising in the same material assuming a slab (disk) geometry and a torus geometry respectively.  In all cases uncertainties have been quite large and further investigation is warranted.   \\subsection{Optical Obscuration and X-ray Absorption}  We have confirmed that \\mcg\\ is unabsorbed in the X-ray band, a fact which is interesting considering its optical classification as a Seyfert 1.5 (Winkler 1992; Winter \\etal 2010).  Standard unification schemes would suggest that it should have less obscuration in the X-ray band than a typical Seyfert 2 but \\textsl{more} than a typical Seyfert 1. However this is not the case, as we have obtained a very low upper limit on column density in excess of the Galactic.  The question then, is whether there is material in the line of sight to the optical emission from this AGN that is \\textsl{not} in the line of sight to the X-ray emitting region.  Optical reddening from dust can be characterized by the visual extinction ($A_\\textrm{v}$), which we calculated from the flux ratio of H$\\alpha$ to H$\\beta$, using the observed values of Winkler (1992) and assuming the intrinsic value of the Balmer decrement to be 2.87 (Osterbrock 1989).  We found $A_\\textrm{v}$\\,$\\sim$\\,1.97~mag. From this value we calculated the inferred column density of gas using the relation of Predehl \\& Schmitt (1995),~ \\NH\\,=\\,$(1.79 \\times 10^{21}$ cm$^{-2})$\\,$A_\\textrm{v}$, and assuming the Galactic gas/dust ratio holds. We inferred a column density of $\\sim$\\,3.55$\\times 10^{21}$ cm$^{-2}$, an order of magnitude higher than our upper limit on the observed X-ray absorption column (in excess of the Galactic) of 2.5$\\times 10^{20}$ cm$^{-2}$. Therefore, assuming the X-ray absorbing gas and optical absorbing dust track each other, the dust obscuring the optical broad line region is most likely \\textsl{not} in the line of sight to the X-ray emitting regions.  The inferred column density is also far too low to be associated with the dusty torus out of the line of sight which must be Compton-thick in order to produce the strong reflection component that we see (\\mytorus gives a column density of 3.6\\err{1.6}{0.9}$\\times 10^{24}$ cm$^{-2}$).  Since the X-ray emitting region is theorized to be very close in to the central black hole, it seems unlikely that this dust is in the form of an extended cloud in the host galaxy lying far (kiloparsecs) from the black hole that just happens to have a hole in the right place to produce a bare nucleus in the X-rays. The dust must be distinct, however, from the Compton-thick torus, given the low inferred column density and evidence that we are viewing the AGN more or less face on. One possible explanation for this set of constraints is that there is clumpy material in the line of sight and it happens that none of the clumps are obscuring the X-ray emitting region (see, e.g., Nenkova \\etal 2008), only the optical broad line region. If future X-ray monitoring discovers a sudden, short-term increase in \\NH, this scenario may be supported. Another explanation is that the material is commensurate with the broad line region, thus obscuring only this region.   \\subsection{Conclusions}  \\suzaku is an ideal tool for studying AGN, providing us with a detailed look at many crucial components of AGN spectra. We confirmed that \\mcg\\ is extremely unabsorbed in the X-ray band (an upper limit of 2.5$\\times 10^{20}$ cm$^{-2}$ in excess of the Galactic column), despite significant reddening seen in the optical band. These results led us to conclude that while no absorbers are in the line of sight to the central region of the AGN, it is possible that clumps of material may be obscuring lines of sight to the region(s) where the optical line emission is produced.  With the excellent resolution of the XIS CCDs around $\\sim$\\,5--7 keV we were able to accurately study the Fe K complex. We found a narrow ($v_\\textrm{FWHM} <$\\,7100 km\\,s$^{-1}$) Fe \\Ka\\ emission line and were able to constrain the location of the emitting material to much farther out than has been done with previous observations. We detected no significant broad Fe line as we would expect to see from the inner portions of a radiatively efficient accretion disk. Since we have such a clean line of sight to the nucleus, this region cannot be simply obscured or out of our line of sight, implying that the inner disk may be truncated or radiatively inefficient.  From our limits on the Fe line $v_\\textrm{FWHM}$ we have calculated a minimum inner radius of $\\gtrsim$\\,1200 $R_{\\rm S}$.  We were also able to associate the Fe emission complex with the Compton reflection component in our models.  Our results indicate that both components likely arise in the same Compton-thick material without any contribution from additional Compton-thin circumnuclear material. Assuming a disk/slab geometry (the \\pexrav model in XSPEC) for the Compton-thick material gave a reflection strength $R$=0.69$\\pm$0.05 with a photon index of $\\Gamma$=1.80$\\pm$0.02.  We successfully applied the new \\mytorus model for Compton reflection, which assumes that the reflecting material is in the form of a torus of uniform density rather than in the form of a flat disk, as has typically been done with more established models such as \\pexrav. Our results gave a photon index of $\\Gamma$=1.70$\\pm$0.01 with a column density of \\NH=3.6\\err{1.3}{0.8}$\\times 10^{24}$ cm$^{-2}$ for the torus. Additionally we found a lower than expected normalization for the Fe line, possibly due to subsolar abundances of Fe with Z$_{\\rm Fe}$=0.75$\\pm$0.14."
    },
    "1107/1107.2075_arXiv.txt": {
        "abstract": "The chemical composition of a protoplanetary disk is determined not only by in situ chemical processes during the disk phase, but also by the history of the gas and dust before it accreted from the natal envelope. In order to understand the disk's chemical composition at the time of planet formation, especially in the midplane, one has to go back in time and retrace the chemistry to the molecular cloud that collapsed to form the disk and the central star. Here we present a new astrochemical model that aims to do just that. The model follows the core collapse and disk formation in two dimensions, which turns out to be a critical upgrade over older collapse models. We predict chemical stratification in the disk due to different physical conditions encountered along different streamlines. We argue that the disk-envelope accretion shock does not play a significant role for the material in the disk at the end of the collapse phase. Finally, our model suggests that complex organic species are formed on the grain surfaces at temperatures of 20 to 40 K, rather than in the gas phase in the $T>100$ K hot corino. ",
        "introduction": "We live on a world rich in complex organic molecules, all of which ultimately derive from a cold pre-stellar core with a much simpler chemical composition. How did the chemistry in the solar system evolve from that original simple state to its current complex state? This chapter focuses on one particular part of the entire evolutionary track: the transition from a pre-stellar core to a protoplanetary disk.  The chemical evolution in a collapsing envelope has been studied extensively with one-dimensional (1D) spherical models (e.g., \\cite[Bergin \\& Langer 1997]{bergin97b}; \\cite[Rodgers \\& Charnley 2003]{rodgers03a}; \\cite[Lee et al. 2004]{lee04a}; \\cite[Aikawa et al. 2008]{aikawa08a}), but such models cannot properly describe the circumstellar disk. We have recently constructed a new collapse model that describes the envelope and the disk in two dimensions (2D) as an axisymmetric structure (\\cite[Visser et al. 2009a]{visser09a}, \\cite[2011]{visser11a}; \\cite[Visser \\& Dullemond 2010]{visser10a}). By taking the vertical extent of the disk into account, we find that accretion from the envelope onto the disk largely takes place near the cold outer edge, rather than near the warmer inner parts as suggested by earlier models. The lower temperatures in our scenario allow for a larger fraction of pristine material to enter the disk. The 2D models also show strong evidence of layered accretion in the disk: new material arriving from the envelope always accretes on top of the older material already in the disk (Fig.\\ \\ref{fig:acclayers}). Another improvement in our model is that we calculate the dust temperatures throughout the disk and the envelope with a full radiative transfer code. The approximate temperature calculations used by some older models are not accurate enough for a reliable investigation of the chemistry.  \\begin{figure}[t] \\begin{center} \\resizebox{0.8\\hsize}{!}{\\includegraphics{visserf1.eps}} \\caption{Schematic view of layered accretion. The outer parts of the original cloud core (left) end up as the surface layers of the circumstellar disk (right). The white part of the core is roughly the part that disappears into the outflow. The arrows illustrate the flow of material onto the disk. If the vertical extent of the disk is not taken into account properly, as happened in some older models, such streamlines would extend further in and the radii at which material accretes would be underestimated.} \\label{fig:acclayers} \\end{center} \\end{figure}  We provide here a simple summary of our model (Sect.\\ \\ref{sec:mod}) and highlight some of the first results (Sect.\\ \\ref{sec:res}). In particular, we focus on the chemical history of H$_2$O and other simple molecules as material flows from the pre-stellar core to the planet-forming disk. The discussion (Sect.\\ \\ref{sec:disc}) addresses two open questions in star formation: how strong is the disk-envelope accretion shock, and do hot corinos really exist? ",
        "conclusions": "We have constructed a two-dimensional axisymmetric model to simulate the collapse of a pre-stellar core into a young star and its surrounding disk. We use the model to study the chemical evolution of the gas and the dust streaming from the envelope to the disk. Comparing the streamlines to what happens in one-dimensional models, we conclude that it is imperative to study core collapse and disk formation in 2D\\@. Our model predicts chemical stratification in the disk due to different physical conditions (density, temperature, radiation field) encountered along different streamlines. We argue that the disk-envelope accretion shock does not play a significant role for the material in the disk at the end of the collapse phase. Finally, our model suggests that complex organic species are formed on the grain surfaces at temperatures of 20 to 40 K, rather than in the gas phase in the $T>100$ K hot corino."
    },
    "1107/1107.2596_arXiv.txt": {
        "abstract": "We report the discovery of Kepler-15b, a new transiting exoplanet detected by NASA's {\\it Kepler} mission. The transit signal with a period of 4.94 days was detected in the quarter 1 (Q1){\\it Kepler} photometry. For the first time, we have used the High-Resolution-Spectrograph (HRS) at the Hobby-Eberly Telescope (HET) to determine the mass of a {\\it Kepler} planet via precise radial velocity (RV) measurements. The 24 HET/HRS radial velocities (RV) and 6 additional measurements from the FIES spectrograph at the Nordic Optical Telescope (NOT) reveal a Doppler signal with the same period and phase as the transit ephemeris. We used one HET/HRS spectrum of Kepler-15 taken without the iodine cell to determine accurate stellar parameters. The host star is a metal-rich ([Fe/H]=$0.36\\pm0.07$) G-type main sequence star with T$_{\\rm eff}=5515\\pm124$~K. The amplitude of the RV-orbit yields a mass of the planet of $0.66\\pm0.1$~M$_{\\rm Jup}$. The planet has a radius of $0.96\\pm0.06$~R$_{\\rm Jup}$ and a mean bulk density of $0.9\\pm0.2$ g\\,cm$^{-3}$. The planetary radius resides on the lower envelope for transiting planets with similar mass and irradiation level. This suggests significant enrichment of the planet with heavy elements. We estimate a heavy element mass of 30-40~M$_{\\oplus}$ within Kepler-15b. ",
        "introduction": "The {\\it Kepler} mission is designed to provide the very first estimate of the frequency of Earth-size planets in the habitable zone of Sun-like stars. The {\\it Kepler} spacecraft continuously monitors $156,453$ stars (Borucki et al.~2011) to search for the signatures of transiting planetary companions. The mission is described in detail in Borucki et al.~(2010). The constant stream of exquisite {\\it Kepler} photometry makes the detection of transits of giant planets relatively easy.  In this paper we describe the first giant planet from the {\\it Kepler} mission confirmed with the Hobby-Eberly Telescope (HET) at McDonald Observatory. Additional RV measurements were also collected with the FIbre--fed \\'Echelle Spectrograph (FIES) at the 2.5\\,m Nordic Optical Telescope (NOT), that are fully consistent with the HET/HRS results. In the following sections we describe the {\\it Kepler} photometry for Kepler-15 and the subsequent ground-based follow-up observations to reject a false-positive and to confirm the planet. Finally we will discuss the radius of Kepler-15b and the planet's internal composition. ",
        "conclusions": "After passing all tests and observational diagnostics, from photometric centroid shifts to spectroscopic line bisectors, and the fact that an RV orbit in period and phase with the transit ephemeris is detected, we conclude that Kepler-15b is indeed a new transiting planet. All our currently available data suggest a planet with a mass of $0.66$~M$_{\\rm Jup}$, a radius of 0.96~R$_{\\rm Jup}$ and a mean density of $0.9\\pm0.2$\\,g\\,cm$^{-3}$ orbiting a metal-rich ([Fe/H]=0.36) G-type star every 4.94\\,days. Kepler-15 is tied with Kepler-6 as the most metal-rich host star of all currently published {\\it Kepler} planets.  Although we initially suspected a large planetary radius, our results demonstrate rather the opposite. Figure\\,\\ref{radius} shows that the planet's radius is modestly smaller than planets of similar mass and irradiation level.  This suggests that the planet is more enriched in heavy elements than most other transiting planets. Given the high stellar metallicity, and the connection between stellar metallicity and planetary heavy elements (Guillot et al.~2006, Burrows et al.~2007, Miller \\& Fortney~2011), this is is not surprising.  Using the tables of Fortney et al.\\,(2007), we estimate a heavy element mass within the planet of at least 30-40 Earth masses."
    },
    "1107/1107.5149_arXiv.txt": {
        "abstract": "To discriminate between  different dust formation processes is a key issue in order to understand its properties. We analysed six submillimeter galaxies at redshifts $4<z<5$ and nine quasars at $5<z<6.4$. We estimated their dust masses from their (sub)millimeter emission and their stellar masses from the spectral energy distribution modelling or from the dynamical and gas masses obtained from the CO line detections. We calculated the dust yields per AGB star and per SN required to explain these dust masses and concluded that AGB stars are not efficient enough to form dust in the majority of these galaxies. SN could be responsible for dust production, but only if dust destruction in the SN shocks is not taken into account. Otherwise even SNe are not efficient enough, which advocates for some other dust production mechanism. We present the hypothesis that grain growth in the interstellar medium is responsible for bulk of the dust mass accumulation in these galaxies. ",
        "introduction": "Dust is of prime importance in cosmology, not only because it obscures our view on stellar populations, but also because its emission contains crucial information about the half of the energy emitted in the Universe. Dust can be formed by asymptotic giant branch (AGB) stars and by supernovae (SNe).  It is known that an AGB star can produce up to $\\sim4\\mbox{}\\times10^{-2}\\,M_\\odot$ of dust \\citep{morgan03,ferrarotti06}, whereas a SN can produce up to $\\sim1.32\\,M_\\odot$ of dust \\citep{todini01,nozawa03}. However, the SN shocks destroy the majority of the dust formed by a SN leaving only $\\lesssim0.1\\,M_\\odot$  to be ejected into the interstellar medium  \\citep[ISM;][]{bianchi07,cherchneff10}.  The dust yields derived observationally for SN remnants are usually $<0.01\\,M_\\odot$ except of Cassiopeia~A \\citep{dunne03,dunne09casA} and Kepler \\citep{morgan03b,gomez09}, for which the derived dust yields are of the order of $\\sim1\\,M_\\odot$.  It is difficult to distinguish the dust formed by AGB stars and by SNe. Based on the flat extinction curves, the SN-origin dust has been claimed to be present in distant quasars (QSOs) and gamma-ray burst host galaxies (\\citealt{maiolino04}, \\citealt{gallerani10}, \\citealt{perley10} and \\citealt{stratta07b}, but see \\citealt{zafar10}).  Another option is that the stellar sources provide only the dust seeds and the bulk of the dust mass accumulation happens in the ISM by the grain growth via capturing of heavy elements  \\citep{draine79,dwek80,draine90,draine09}.  The significant amounts of dust have been detected in the early Universe up to $z\\sim6.4$ \\citep{dunlop94,benford99,archibald01,omont01,priddey01,priddey03,priddey08,isaak02,bertoldi03,robson04,beelen06,wang08b,martinezsansigre09,michalowski10smg,santini10}. Hence, its origin has to be explained by a process that is both efficient and rapid.  Several works have attempted modelling of the dust evolution at these high redshifts  \\citep{dwek07,valiante09,dwek11,gall11,gall11b,pipino11}. In this paper we present more simplistic approach of estimating whether stellar sources are efficient enough to explain the dust detected at high redshifts. ",
        "conclusions": "We investigated the possible dust producers in the SMGs at $4<z<5$ and QSOs at $5<z<6.4$ by estimating the required dust yields per AGB star and per SN and comparing them to the allowed values  derived from the theoretical models. We found that AGB stars are not efficient to form dust in the majority of these galaxies leaving SNe as the only likely stellar dust producers. However, the required dust yields for SNe exceed the allowed values if dust destruction in the SN shocks is taken into account. This advocates for some other dust production mechanism e.g.~the grain growth in the ISM. This process is fast enough and SNe deliver enough heavy elements to build up the dust masses derived for these galaxies."
    },
    "1107/1107.4433_arXiv.txt": {
        "abstract": "Motivated by the dawn of precision cosmology and the wealth of forthcoming high precision and volume galaxy surveys, in this paper we study the effects of inhomogeneities on light propagation in a flat $\\Lambda$CDM background. To this end we use exact solutions of Einstein's equations \\citep{MeuBru11-1} where, starting from small fluctuations, inhomogeneities arise from a standard growing mode and become non-linear. While the matter distribution in these models is necessarily idealised, there is still enough freedom to assume an arbitrary initial density profile along the line of sight. We can therefore model over-densities and voids of various sizes and distributions, e.g. single harmonic sinusoidal modes, coupled modes, and more general distributions in a $\\Lambda$CDM background. Our models allow for an exact treatment of the light propagation problem, so that the results are unaffected by approximations and unambiguous.  Along lines of sight with density inhomogeneities which average out on scales less than the Hubble radius, we find the distance redshift relation to diverge negligibly from the Friedmann-Lemaitre-Robertson-Walker (FLRW) result. On the contrary, if we observe along lines of sight which do not have the same average density as the background, we find large deviations from the FLRW distance redshift relation. Hence, a possibly large systematic might be introduced into the analysis of cosmological observations, e.g.\\ supernovae, if we observe along lines of sight which are typically more or less dense than the average density of the Universe. In turn, this could lead to wrong parameter estimation: even if the Cosmological Principle is valid, the identification of the true FLRW background in an inhomogeneous universe maybe more difficult than usually assumed. ",
        "introduction": "At the very basis of modern cosmology lies the assumption the Universe is, at any given time, homogeneous and isotropic on large scales. This is translated mathematically into a Robertson-Walker metric, i.e. a metric that is assumed to represent a space average and is therefore exactly homogeneous and isotropic. In addition, the non-trivial hypothesis is made that this metric should be a solution of Einstein's equations, thereby giving rise to a FLRW universe model. Having assumed the Cosmological Principle, the growth of inhomogeneities and their effects are typically modelled with perturbation theory about a ``background'' FLRW model. Within this framework, the formation of non-linear structures at smaller scales is considered in N-body simulations using the Newtonian approximation \\citep{Spr05}. Most observations are interpreted assuming this Friedmannian framework; in particular, distances are computed assuming a FLRW distance-redshift relation, i.e. completely neglecting inhomogeneities.  We are currently living in a time when galaxy surveys and other observations are reaching unprecedented sky coverage and precision; therefore, it seems timely to fully investigate the effects of the non-linear growth of structures on observations, within a general relativistic framework. Much work has been done in trying to understand the effect of inhomogeneities on observations by using perturbations around an FLRW model \\citep{DyeRoe72, DyeRoe73, DyeRoe74, Sas87, FutSas89, PynBir04, BarMatRio05, BonDurGas06, Ras11-2}. However, these analyses are limited to linear structure growth and therefore cannot properly take into account non-linear inhomogeneities. Non-linearities can be modelled using idealised matter distributions. Several different approaches using generalised Swiss-Cheese models or alternative geometries have been considered in e.g.\\ \\citet{BroTetTza07, Mar07, BroTetTza08, BisNot08, CliZun09, BolCel10, Szy11, NwaIshTho11}, where significant deviations are usually only found for very large scale density inhomogeneities. On the contrary, claims are made that small scale structure formation might have a back-reaction effect on the overall expansion of the Universe, see \\citet{Ras06} and \\citet{Buc08} for overviews of this topic. We do not investigate this phenomenon of back-reaction in this paper though, as our model clearly splits into inhomogeneities and background dynamics.  A class of models which are very appropriate for considering a discrete distribution of matter in an otherwise FLRW expanding universe has recently been analysed in depth by \\citet{CliFer09, Clifer09-2}. In these models, originally introduced by \\citet{LinWhe57} and revised by \\citet{Red88}, matter is described by pointlike masses in a spherically symmetric void box (represented by Schwazschild space-times) and these boxes are distributed in a lattice and the overall expansion is described by the Friedmann equation. The main motivation of Clifton and Ferreira in following the Lindquist and Wheeler construction is the observation that the Universe largely consists of galaxies and clusters of galaxies surrounded by vacuum. The question they address is how observations and measurements of the cosmological parameters are affected in a highly inhomogeneous universe whose overall dynamics are homogeneous and isotropic. However, this lattice construction is only an approximate solution to Einstein's equations and has regions of ``no man's land\" in between the matched spheres which might have an effect on the light tracing, see \\citet{Cli10}. Perhaps even more importantly, the inhomogeneities in this model are strongly non-linear at all times.  Having the same type of questions addressed by \\citet{CliFer09} in mind, in this paper we shall investigate the optical properties of an exact solution to Einstein's field equations (EFEs), developed in our previous paper \\citep{MeuBru11-1}. This solution belongs to the second class Szekeres models \\citep{Sze75} for dust, generalised by \\citet{BarSte84} to include the cosmological constant $\\Lambda$. In \\citet{MeuBru11-1} we used the \\citet{GooWai82} formulation of the Szekeres models to reconsider the Barrow and Stein-Schabes solutions in a $\\Lambda$CDM context. In our model, starting from standard small perturbations of a $\\Lambda$CDM homogeneous and isotropic universe, the matter distribution is continuous and can evolve to a highly non-linear stage. In the process, the inhomogeneities can either form a distribution of large voids, over-densities, or a mixture of the two with over-densities possibly even forming pancakes as in the Zel'dovich approximation in Newtonian cosmology. The benefit of our model is therefore two-fold: \\textit{i}) we consider exact solutions of Einstein's equations, therefore avoiding any possible problem associated with approximations and matching and \\textit{ii}) these exact solutions describe non-linear inhomogeneities growing on top of a FLRW background with the possibility of modelling a rather arbitrary distribution of both voids and over-densities.  In particular, we shall consider the effect of inhomogeneities on the redshift, angular diameter distance and distance modulus. We begin by considering single mode harmonic sinusoidal deviations from homogeneity and then the case of coupled modes. Finally, closer in spirit to the \\citet{CliFer09} work, we consider inhomogeneities where peaks and voids, arranged in a periodic array, are more strongly separated than can be achieved by simple harmonic distributions. We demonstrate that the deviations from the FLRW background in the determination of the distances is mainly due to the Ricci and Weyl focusing terms in the Sachs equations and show that instead the shear of the null congruence has a negligible effect. We also briefly investigate the effect of mode coupling on the growth of structure and, interestingly, we show that even a long wavelength mode with small amplitude can strongly enhance the growth of short wavelength modes, thanks to the non-linear coupling. The non-linear interaction of different modes does not seem to influence the distance measures significantly and we find that changes in the redshift and distances are mostly affected by the long wavelength modes. Using an array of density profiles which are not sinusoidal but quite peaked around the maximum and separated by large voids, we find that the effect on the redshift and distance measures does not prove to be significantly more than using an initially sinusoidal density distribution with the same wavelength of the array scale. Overall, we find that all deviations in the redshift and distance measures are less than 1\\%, when we consider what we refer to as ``compensated inhomogeneities\" along the line of sight, i.e.\\ where the average density along the line of sight matches the background density. However, this does not need to be the case: when the inhomogeneities are on average above or below the background, the effects on redshift and distance measures can be very large.  A summary of the paper is as follows. In Sec.\\ \\ref{sec:sol} we will briefly present the exact solution we will be using in this work, referring the reader to \\citet{MeuBru11-1} for more details. Subsequently, in Sec.\\ \\ref{sec:geo}, we will derive the null geodesic equations for the given metric and in Sec.\\ \\ref{sec:opt} we derive the form that the Sachs optical equation will take for a given physical situation that is to be investigated. In Sec.\\ \\ref{sec:res} we shall present the results of our analysis, considering single mode deviations in Sec.\\ \\ref{sec:res1}; multiple modes and their coupling and effects in Sec.\\ \\ref{sec:res2}; and the effects of an array of strong peaks and large voids in Sec.\\ \\ref{sec:res3}. In Sec.\\ \\ref{sec:conc} we draw our conclusions. In Appendix \\ref{sec:trans} we present details on the tetrad transformations needed to derive the Sachs optical equations in our model.  Throughout the paper we choose units $c=8\\pi G=1$ and assume the standard $\\Omega_{\\Lambda}=0.75$ and in the commonly used units $H_0=72\\text{km}\\text{s}^{-1}\\text{Mpc}^{-1}$. ",
        "conclusions": "\\label{sec:conc} In this paper we have analysed the effects of non-linear structure on redshift and distance measures using the exact solution developed in our previous paper \\citep{MeuBru11-1}. This model, described by the line element (\\ref{eq:line}) in synchronous comoving coordinates, allows us to choose an arbitrary matter distribution along one line of sight with the growth rate of structure and the density distribution away from this axis being set by EFEs. A remarkable feature of our model is that the inhomogeneities are described by a single metric function $F$, which satisfies the same linear second order differential equation satisfied by the linear density perturbation $\\delta$ in Newtonian perturbation theory as well as in relativistic perturbations when the synchronous comoving gauge is used. Therefore $F$ satisfies a superposition principle and, in particular, extends into the non-linear regime the same growing and decaying modes that $\\delta$ shows in the linear regime.  We have developed the null geodesic equations and the Sachs optical equations in our model for light rays travelling in arbitrary directions. This set-up has then been applied to different physical situations, considering single sinusoidal mode deviations in the density, the coupling of two and three harmonic modes as well as a more complex matter distribution described by an array of peaks and voids along the line of sight. Furthermore, we have investigated which terms in the Sachs optical equations are mainly responsible for deviations from the FLRW values. Additionally, we have analysed the interaction of two and three modes in the growth of structure within our exact non-linear framework.  We consider the redshift and distance measures for single mode density deviations on different length scales and of different amplitudes. The largest effect for the redshift and distance measure to deviate from the FLRW results seems to be obtained for larger density deviations and larger scales. The results are shown in Figs.\\ \\ref{fig:dev500}, \\ref{fig:dev100} and \\ref{fig:dev1}. Given our conservative assumptions on the density profiles, all deviations are below the $1\\%$ level.  Even if the metric function $F$ satisfies a linear equation, our model is non-linear and so the effects of two modes on the redshift and distances does not need to be the same as their combined effects, hence we considered how modes interact in the growth of the density deviations and how the redshift and distance measures are affected by them. We find that combining small and large scale density deviations has significant effects on the growth of structure, in that the peaks of the large scale modes significantly enhance the growth of small scale deviations, see Fig.\\ \\ref{fig:deltas} for plots illustrating this point and the text in Sec.\\ \\ref{sec:res2}. For more details on the growth of structure in our model, see \\citet{MeuBru11-1} and for a more general discussion of the relativistic behaviour of inhomogeneities and overall properties of the class of solutions we consider, see \\citet{MatPanSae94-1,MatPanSae94,BruMatPan95-2,BruMatPan95-1}. The interaction between short and long wavelength modes observed in our model does not seem to have a significant effect on the redshift and distances measures though, as the effect of the long wavelength deviation remains dominant despite the presence of small scale deviations, see Fig.\\ \\ref{fig:dev1100} for the results of a $1$ Mpc and $100$ Mpc (today) mode combination. This implies that mode coupling does not provide significantly larger deviations on the redshift and distances than the individual modes.  To generalise our result to density distributions where peaks and voids are more pronounced than in a single mode sinusoidal, we have considered an array of density profiles given in Eq.\\ (\\ref{eq:cosh}), which provides quite peaked over-densities and large voids separating them, choosing a typical array scale of $100$ Mpc today. The results of the light tracing for this distribution is given in Fig.\\ \\ref{fig:cosh}, where the density distribution is shown in the top panel. Given this density profile, which is clearly non-sinusoidal for all times, the deviations from the FLRW redshift and distance measure are still comparable in amplitude to the results found in the single $100$ Mpc mode analysis, with the large scale deviations being below the $1\\%$ level.  For all these different density distributions, we investigated which of the terms in the Sachs optical equations is dominant in providing the deviations, see Fig.\\ \\ref{fig:focus} for the different terms for a single mode deviation. We find that in all cases the Ricci focusing term is dominating, while the effect of the shear on the angular diameter distance seems to be vanishingly small. Due to its special geometric character (the space-time is Petrov type D), in our model the Weyl focusing is exactly zero along the $z$-axis of symmetry and it is sub-dominant with respect to the Ricci focusing in direction at an angle $\\alpha$ with respect to the $z$-axis because of a $\\sin^2(\\alpha)$ factor. However, both the Ricci and Weyl focusing are proportional to the density deviation $\\delta$ and so we may expect that in a more general space-time they would be of the same order.  All the above mentioned results are for density deviations which we refer to as compensated, i.e. where the metric function $F$ averages to zero along the line of sight at all times, so that the initially small density profile is also compensated. The results from considering density profiles which do not average to zero initially, however, are quite different. In Fig.\\ \\ref{fig:dzda}, we show that over-dense and under-dense lines of sight have significant effects on the redshift and angular diameter distance. This implies that if we, on average, observe along lines of sight which are more or less dense than the background, we may need to expect significant effects. Therefore it is important to understand whether the lines of sight we observe along are really average ``skewers\" through the Universe matter distribution. It also emphasises the importance of identifying the correct background density with observations, as an over- or under-estimate may affect our interpretation of observational results significantly. Understanding these effects is crucial to our interpretation of cosmological observations and hence we leave a deeper analysis for future work.  Finally, part of the motivation for our work has come from the strong deviations from the standard FLRW results in \\citet{CliFer09}, and therefore we would like to briefly compare our results here. Considering density distributions which are initially compensated along the line of sight cannot provide deviations in the redshift and distance measures as large as found in \\citet{CliFer09}. This does seem to be in agreement with most Swiss-Cheese type analyses. However, considering lines of sight which have an average density which is lower than the background density can provide similar results as found in \\citet{CliFer09}."
    },
    "1107/1107.4119_arXiv.txt": {
        "abstract": "Galaxy clusters grow by mergers with other clusters and galaxy groups. These mergers create shocks within the intracluster medium (ICM). It is proposed that within the shocks particles can be accelerated to extreme energies. In the presence of a magnetic field these particles should then form large regions emitting synchrotron radiation, creating so-called radio relics. An example of a cluster with relics is CIZA~J2242.8+5301. Here we present hydrodynamical simulations of idealized binary cluster collisions with the aim of  constraining the merger scenario for this cluster. We conclude that by using the location, size and width of double radio relics we can set constraints on the mass ratios, impact parameters, timescales, and viewing geometries of binary cluster merger events. ",
        "introduction": "\\label{sec:intro} Radio relics are elongated filamentary sources found in massive merging galaxy clusters. These radio sources indicate the presence of magnetic fields and particle acceleration within the ICM \\citep[e.g.,][]{1977ApJ...212....1J, 2004IJMPD..13.1549G}. Since all giant radio relics are found in merging clusters \\citep[e.g.,][]{2004ApJ...605..695G, 2007A&A...467...37B, 2009A&A...507..661B, 2010ApJ...721L..82C}, it has been proposed that a small fraction of the energy released during a cluster merger event is channeled into (re)acceleration of particles. The currently favored scenario for the origin of giant radio relics is that they trace shock waves within the ICM in which particles are (re)accelerated by the diffusive shock acceleration mechanism \\citep[e.g.,][]{1983RPPh...46..973D, 1987PhR...154....1B, 1991SSRv...58..259J, 2001RPPh...64..429M}. An alternative scenario has been proposed by \\cite{2010arXiv1011.0729K} based on a single secondary cosmic ray electron model, where the time evolution of both magnetic fields and cosmic ray distribution are taken into account to explain giant relics (and radio halos).  Particularly interesting are the so-called double-relics, where two relics are diametrically located on both sides of the cluster center \\citep[e.g., ][]{2009A&A...494..429B, 2009A&A...506.1083V, 2006Sci...314..791B, 2010Sci...330..347V,2011ApJ...727L..25B,2011A&A...528A..38V}. These relics are thought to trace outward moving shock waves from a binary cluster merger event, which develop after core passage of the two subclusters. The idea is that double relics can be used to put constraints on the merger timescale, mass ratio, impact parameter and viewing angle. ",
        "conclusions": "\\label{sec:summary} We have  simulated radio relics  from idealized binary cluster merger events, varying the mass ratios, impact parameters and viewing angles. The radio emission in these simulations is produced by relativistic electrons which are accelerated via the DSA mechanism.  We will use these simulated radio maps to constrain the merger scenario for the cluster CIZA~J2242.8+5301, which hosts a Mpc scale double radio relic. A preliminary analysis indicates that a merger event with a mass ratio of $1.5:1-2.5:1$ provides the best match to the observed radio emission in this cluster at 1382~MHz. The impact parameter of the merger is constrained to be $b \\lesssim 400$~kpc, while the shock fronts (and relics) are observed under and angle less than $\\sim10\\deg$ from edge-on. The relics are seen at a time of about 1~Gyr after core passage. Our simulations suggest that double radio relics provide a powerful tool to determine the merger parameters when X-ray observations are not available. One of the main uncertainties is the structure and strength of magnetic field in the ICM. Magneto-hydrodynamical simulations of cluster mergers will be needed to study the evolution of the magnetic field and its influence on the radio emission from merger shocks."
    },
    "1107/1107.3739_arXiv.txt": {
        "abstract": "We consider inflation within the context of what is arguably the simplest non-metric extension of Einstein gravity. There non-metricity is described by a single graviscalar field with a non-minimal kinetic coupling to the inflaton field  $\\Psi$, parameterized by a single parameter $\\gamma$. We discuss the implications of non-metricity for chaotic inflation and find that it significantly alters the inflaton dynamics for field values $\\Psi \\gtrsim M_P/\\gamma$, dramatically changing the qualitative behaviour in this regime. For potentials with a positive slope non-metricity imposes an upper bound on the possible number of e-folds. For chaotic inflation with a monomial potential, the spectral index and the tensor-to-scalar ratio receive small corrections dependent on the non-metricity parameter. We also argue that significant post-inflationary non-metricity may be generated. ",
        "introduction": "Even at the classical level, the theory of gravity has still two major open issues. First, what is the form of the action? Here investigations have focused on extensions of Einstein gravity, such as $f(R)$ gravities \\cite{Sotiriou:2008rp}, and their implications for dark energy or inflation \\cite{Nojiri:2006ri,Clifton:2011jh}. But there is also the second, and arguably even more profound question: what are the true gravitational degrees of freedom? Conventionally, one just assumes metricity: that gravity can be described by the metric alone, setting the connection $\\hat\\Gamma$ to be the Levi-Civita connection with $\\hat\\Gamma=\\Gamma[g_{\\mu\\nu}]$ by hand. In the Palatini variation one takes the metric and the connection to be independent degrees of freedom \\cite{Hehl:1994ue}, and varying both results in equations of motion that differ from the metric case except if the Lagrangian has exactly the Einstein-Hilbert form with $f(R)=R$. In the general Palatini case, the solution to equation of motion for the connection is not the Levi-Civita connection, and hence such theory is called non-metric. In non-metric gravity, the metric is no longer covariantly conserved.  However, several issues arise in the Palatini approach \\cite{Olmo:2011uz}, which may stem from the nontensorial property of the connection. Instead, one could consider theories with two metrics: one for describing the geometry of the manifold, and the other for generating the connection. Indeed, the Palatini case can be seen as a special case of a more general class of conformal non-metric theories. In these C-theories \\cite{Amendola:2010bk,Koivisto:2011vq} one postulates a metric for the connection, $\\hat g_{\\mu\\nu}$ that is related to the gravitational metric by a conformal factor $\\C$ with $\\hat g_{\\mu\\nu}=\\C g_{\\mu\\nu}$, the Palatini-f(R) then corresponding to the choice $\\C(R)=f'(R)$. The connection is then the Levi-Civita connection of the connection-generating metric with $\\hat\\Gamma=\\Gamma[\\hat g_{\\mu\\nu}]$. Such a relation was originally considered by Weyl in his conformally invariant theory of gravity \\cite{0264-9381-13-1-013}.  Surprisingly, the subtle refinement of the fundamental spacetime structure turns out to have concrete and interesting consequences, for one allowing non-metric degrees of freedom to propagate in vacuum. This is in contrast both to the result of the standard metric variation (where the connection is fixed to be metric a priori) and the metric-affine, i.e. Palatini, variation (where the connection is constrained a posteriori), which however can be recovered as specific limits of C-theories. Thus, they provide a unified approach to a variety of previously seemingly very different alternative gravity theories but reveal also novel possibilities hidden already in the Einstein-Hilbert action itself.    The nature of gravity is of great relevance for models of inflation \\cite{Mazumdar:2010sa}. % Since inflation occurs at very high curvature scales where one naturally expects corrections to Einstein's theory, the ambiguities both in the form of gravitational Lagrangian and its degrees of freedom should be taken into account. The purpose of this paper is to explore the implications of the simplest conformal non-metricity to inflationary physics. In particular, we analyze the corrections that inevitably appear to large-field chaotic inflation. Our approach is different from considering inflation to be driven by non-metricity, which has been pursued in Refs. \\cite{0264-9381-5-3-013,0264-9381-8-5-014}. The Palatini variation has been applied to a non-minimally coupled inflation \\cite{Bauer:2010jg,Tamanini:2010uq}, however there taking only into account the shift of the kinetic term \\cite{Koivisto:2005yc}, whereas our purpose is to consider the generic lagrangian (for details, see appendix \\ref{quadratic}).  In the next section, we couple a matter scalar field into gravity in the presence of Weyl non-metricity and derive the equations of motion for homogeneous fields. In the following section we then analyze the dynamics of slow-roll chaotic inflation and derive constraints on the amount of non-metricity chaotic inflation can tolerate. The results are discussed in section \\ref{conclusions}. The generic starting lagrangian and the general field equations are confined to the appendices. ",
        "conclusions": "\\label{conclusions} The simplest possible implementation of a non-metric theory of gravity introduces a new scalar degree of freedom associated with the conformal part of the metric. In turn, this can be cast as a theory of two scalar fields with a non-trivial metric in field space. We have examined the implications of such a non-metric theory for inflation and found that for inflaton values $\\Psi\\gtrsim 1/\\gamma\\kappa$ such a seemingly innocuous modification has profound implications for the way inflation proceeds. In particular, as can be seen from eq (\\ref{u1}) the effect of non-metricity is to effectively turn an attractive inflaton potential into a repulsive one and viceversa. Thus, an inflationary potential with a positive slope becomes unstable in the region $\\Psi\\gtrsim 1/\\gamma\\kappa$ leading to a cosmology that cannot represent the inflationary past of our universe. On the contrary, for values $\\Psi < 1/\\gamma\\kappa$ slow roll inflation can proceed with few modifications, mainly in the form of small $\\gamma$ corrections for the spectral index and the scalar to tensor ratio. The main difference from metric inflation is that now the non-metricity parameter $\\gamma$ sets un upper bound for the possible number of e-folds, see eq (\\ref{efolds1}). In this sense non-metricity can act as an IR regulator for inflation.  It is perhaps interesting to note that potentials which are repulsive can turn attractive in the $\\Psi\\gtrsim 1/\\gamma\\kappa$ region which could then support a viable inflationary cosmology, although again with a point of infinite force for the field at $\\Psi= 1/\\gamma\\kappa$.   Let us close by briefly discussing the effects of perturbations. We argued that non-metric inflation would give rise to a fluctuation of the graviscalar $\\varphi$ in addition to the standard adiabatic metric perturbation, see eq (\\ref{isocurv}). After the inflaton decays the graviscalar perturbation would presumably persist along with the adiabatic metric one $\\zeta(k)\\simeq{H_*}/{\\sqrt{\\epsilon}k^{3/2}}$. This would give rise to a non-metric contribution to the perturbed part of the connection, see equation (\\ref{conformal}), with $(\\hat{\\Gamma}-\\Gamma)/\\Gamma\\sim\\sqrt{\\epsilon}$. Thus, non-metric inflation seems to be generating deviations of the perturbed post-inflationary connection from the metric case. Such deviations could be used to constrain non-metricity further and we will return to this issue in future work.   \\subsubsection*{Acknowledgements}  \\noindent  KE is supported by the Academy of Finland grants 218322 and 131454. GR is supported by the Gottfried Wilhelm Leibniz programme of the Deutsche Forschungsgemeinschaft.                  \\appendix"
    },
    "1107/1107.1957_arXiv.txt": {
        "abstract": "{In this paper I describe a simple numerical procedure to compute synthetic horizon altitude profiles for any given site. The method makes use of a simplified model of local Earth's curvature, and it is based on the availability of digital elevation models describing the topography of the area surrounding the site under study. Examples constructed using the Shuttle Radar Topographic Mission (SRTM) data (with 90m horizontal resolution) are illustrated, and compared to direct theodolite measurements. The proposed method appears to be reliable and applicable in all cases when the distance to the local horizon is larger than $\\sim$10 km, yielding a rms accuracy of $\\sim$0.1 degrees (both in azimuth and elevation).  Higher accuracies can be achieved with higher resolution digital elevation models, like those produced by many modern national geodetic surveys.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.4097.txt": {
        "abstract": "We obtain cosmological constraints from a measurement of the spherically averaged power spectrum of the distribution of about 90000 luminous red galaxies (LRGs) across 7646 deg$^{2}$ in the Northern Galactic Cap from the seventh data release of the Sloan Digital Sky Survey. The errors and mode correlations are estimated thanks to the 160 LasDamas mock catalogues, created in order to simulate the same galaxies and to have the same selection as the data. We apply a model, that can accurately describe the full shape of the power spectrum with the use of a small number of free parameters. Using the LRG power spectrum, in combination with the latest measurement of the temperature and polarisation anisotropy in the cosmic microwave background (CMB), the luminosity-distance relation from the largest available type 1a supernovae (SNIa) dataset and a precise determination of the local Hubble parameter, we obtain cosmological constraints for five different parameter spaces. When all the four experiments are combined, the flat $\\Lambda$CDM model is characterised by $\\Omega_{\\mathrm{M}} = 0.259_{-0.015}^{+0.016}$, $\\Omega_{\\mathrm{b}} = 0.045 \\pm 0.001$, $n_{\\mathrm{s}}=0.963\\pm0.011$, $\\sigma_8 = 0.802\\pm0.021$ and $H_0 = 71.2\\pm1.4\\,\\mathrm{km}\\,\\mathrm{s^{-1}\\,Mpc^{-1}}$. When we consider curvature as a free parameter, we do not detect deviations from flatness: $\\Omega_{\\mathrm{k}} = \\left(1.6\\pm5.4\\right)\\times10^{-3}$, when only CMB and the LRG power spectrum are used; the inclusion of the other two experiments do not improve substantially this result. We also test for possible deviations from the cosmological constant paradigm. Considering the dark energy equation of state parameter $w_{\\mathrm{DE}}$ as time independent, we measure $w_{\\mathrm{DE}} = -1.025_{-0.065}^{+0.066}$, if the geometry is assumed to be flat, $w_{\\mathrm{DE}} = -0.981_{-0.084}^{+0.083}$ otherwise. When describing $w_{\\mathrm{DE}}$ through a simple linear function of the scale factor, our results do not evidence any time evolution. In the next few years new experiments will allow to measure the clustering of galaxies with a precision much higher than achievable today. Models like the one used here will be a valuable tool in order to achieve the full potentials of the observations and obtain unbiased constraints on the cosmological parameters. ",
        "introduction": "The last decade was characterised by a dramatic change in our vision of the Universe and by an impressive increase in size and quality of available datasets. At the end of the twentieth century, the analysis of the luminosity-distance relation in the Type 1a supernovae (SNIa) showed that the Universe is undergoing a phase of accelerated expansion driven by a new exotic component, dubbed dark energy, whose energy density is about 70\\% of the total \\citep{riess_98,perlmutter_99}. This has been then confirmed by other observations, such as the cosmic microwave background \\citep[CMB, e.g.,][]{hinshaw_wmpa1_aps, Spergel_03, Spergel_07, komatsu_wmap09, komatsu_10}, the large scale structure of the Universe \\citep[LSS; e.g.][]{efstathiou02, percival02, tegmark04, Sanchez_06, Sanchez_09, percival_10_SDSS, reid_10_SDSS, blake_11} and the number density of galaxy clusters as function of their mass \\citep[e.g.][]{vikhlinin_09}. Different combinations of these probes have been used in the past years to constrain the dark energy equation of state parameter $w_{\\mathrm{DE}}$ with about a 5-10\\% error. The analysis presented here aims at constraining cosmological parameters and shed light on the nature of dark energy.  Many present day and future galaxy redshift surveys, like the Baryonic Oscillation Spectroscopic Survey \\citep[BOSS,][]{schlegel_BOSS, eisenstein_11_sdss3}, the Panoramic Survey Telescope \\& Rapid Response System \\citep[Pan-STARRS,][]{panstarrs}, the Dark Energy Survey \\citep[DES,][]{des}, the Hobby Eberly Telescope Dark Energy Experiment \\citep[HETDEX,][]{hetdex} and the space based Euclid mission \\citep{laureijis_09_euclid}, are designed to constrain the properties of dark energy, usually parametrised by its density, the present day value of $w_{\\mathrm{DE}}$ and its possible time evolution, with unprecedented precision. This will help to reduce the number of possible models of dark energy that has been proposed in the last decade, that are either based on a time varying field \\citep[for a review see, e.g.,][]{Peebles2003}, or on modifications of the equation of general relativity \\citep[modified gravity, for a review see, e.g.,][]{tsujikawa_11}.  In this work we concentrate on the distribution of galaxies on large scales as a means to constrain the history and composition of the Universe, analysing it statistically through the power spectrum, the Fourier transform of the two point correlation function. The shape of the galaxy power spectrum and correlation function contains information about the composition of the Universe and the non-linear evolution of clustering, bias \\citep[e.g.][]{mcdonald_renorm, mats_LPT2, Jeong_09} and redshift space distortions \\citep[e.g.][]{scoccimarro_04, cabre_09a, cabre_09b, Jennings_2010, reid_11_rsdist}. Biasing is due to the fact that the objects that we observe do not trace perfectly the underlying dark matter distribution and redshift space distortions are introduced when inferring the distance of a galaxy from the measured redshift, which is a sum of a cosmological component and a Doppler shift due to the peculiar motion of the emitter. The two-point statistics contain also a feature, called baryonic acoustic oscillations (BAO), that has been advocated as a powerful tool to probe the curvature of the Universe and that has been intensively studied over the past years. BAOs are the relic signature in the matter distribution of the acoustic oscillations in the baryon-photon plasma in the hot young Universe. The BAOs show up in the correlation function as a quasi gaussian bump at scales $r\\approx 100-110 \\Moh$ \\citep{Matsubara_04} and as a sequence of damped quasi-harmonic oscillations at wave-numbers $0.01\\hoM<k<0.4\\hoM$ in the power spectrum \\citep{Sugiyama_95, Eisenstein_98, EH99}. BAOs in the CMB where predicted by \\citet{peebles_70} and \\citet{sunyaev_70} and first detected in the galaxy distribution by \\citet{cole_05_2dF} and \\citet{Eisenstein_05}. The BAO scale, as measured from the CMB, depends only on the plasma physics prior to recombination, which allows in principle to use it as a standard ruler. Non linear evolution damps and shifts by few percent the acoustic peaks \\citep{crocce_nonlinBAO, Sanchez_08, smith_08}. The full shape of the power spectrum and of the correlation function, as well as the BAOs alone, have been used, alone and in combination with other independent datasets, to constrain cosmological parameters  \\citep[e.g.,][]{percival07, sanchez_cole, cabre_09a, gaztanaga09, Sanchez_09, kazin_10a, percival_10_SDSS, reid_10_SDSS, blake_11, tinker_11}.  In order to obtain unbiased constraints from the information encoded in the large scale structure of the Universe, non linear distortions, bias and redshift space distortions need to be accounted for. Perturbation theory \\citep[PT, see][for a review]{Bernardeau_02} can successfully model the shape of the power spectrum and of the correlation function at redshifts larger than $z=1$ or at large scales when including at least third order terms in the density fluctuations \\citep{Jeong_06,Jeong_09}. In the past few years many groups have proposed different improvements over perturbation theory \\citep[e.g.,][]{crocce_RPT1, crocce_RPT2, mcdonald_renormbias, matarrese_rengroup1, bernardeau_08,matarrese_rengroup2, mats_LPT1, mats_LPT2, pietroni_flowtime, taruya_closure, smith09_SZ, taruya_09, Elia_10}. While some of this approaches, like for instance \\citet{mats_LPT2} and \\citet{taruya_10}, can account for bias and redshift space distortions, most of them describe only the clustering of dark matter in real space. For this reason phenomenological approaches based on a given flavour of PT have been put forward in order to model the dark matter halo and galaxy distributions in real and in redshift space \\citep[e.g.,][the latter will be thereafter referred to as M10]{mcdonald_renorm, crocce_nonlinBAO, Sanchez_08, Jeong_09, montesano_10a}. The model proposed by M10 is inspired by renormalised perturbation theory \\citep[RPT][]{crocce_RPT1, crocce_RPT2} and can describe the full shape of the mildly non-linear the dark matter and halo-power spectrum from numerical N-body simulation and allows to obtain unbiased constraints on the dark energy equation of state parameter.  In this analysis we apply the model developed and tested in M10 to the power spectrum measured from the luminous red galaxies (LRGs), released publicly in the seventh data release of the Sloan Digital Sky Survey (SDSS DR7). Combining this measurement with the latest results from CMB, SNIa and the precise determination of the local Hubble constant, $H_{0}$, we obtain cosmological constraints for five different parameter spaces. The LRG sample and the mock catalogues used to estimate the covariance matrix of the data are presented in Section \\ref{sec:lrg_mocks}. The power spectra and covariance matrix computed from the real and mock datasets are shown in Section \\ref{sec:power_spectra}. The CMB and SNIa datasets and the $H_{0}$ measurement are described in Section \\ref{sec:extra_exp}. Section \\ref{ssec:model} gives a short description of the model that we apply to the LRG power spectrum. In Sections \\ref{ssec:parspace} and \\ref{ssec:issues} we illustrate the parameter spaces that we explore in our analysis and the technique used to perform the fits. In Section \\ref{ssec:test} we test our model against the mean power spectrum of the mock catalogues and of the LRG, in order to determine whether it can describe accurately the full shape of the two-point statistic when a complex geometry is used. Our main results, the cosmological constraints for the five parameter spaces, are discussed in Section \\ref{sec:cosmopars} and then compared with some of the most recent works on the topic in Section \\ref{sec:comparison}. Finally we summarise our results and draw our conclusions in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have computed the power spectrum of the distribution of about 90,000 luminous red galaxies, extracted from the spectroscopic part of the seventh data release of the Sloan Digital Sky survey. To compute the covariance matrix for the power spectrum we make use of the 160 LasDamas mock catalogues, which have the same angular and redshift distribution as the observed LRGs. We describe the measured power spectrum with a model inspired by renormalised perturbation theory and modified in order to describe biased objects in redshift space (section \\ref{ssec:model}). This model has been successfully tested against numerical simulations (M10) and mock catalogues (section \\ref{ssec:test}) for $k\\lesssim0.15\\hoM$ at z=0-0.5.  Combining the large scale structure information with measurements of the CMB temperature and polarisation power spectrum from the seven year data release of WMAP, ACBAR, BOOMERanG, CBI and QUAD and with the luminosity-distance relation from the Union2 supernovae sample and using a precise determination of the local Hubble parameter as a gaussian prior, we explore the constraints in five different cosmological parameter spaces described in section \\ref{ssec:parspace}. They are the flat $\\Lambda$CDM concordance model, a similar one where curvature is a free parameter (k$\\Lambda$CDM) and three models in which the dark energy equation of state has a parametric form: in two cases it is assumed to be constant, with and without the assumption of a flat geometry (wCDM and kwCDM), and in the last case (waCDM) we model $w_{\\mathrm{DE}}$ with the simple parametric formula of equation (\\ref{eq:wa}).  Overall, we obtain tight constraints on the cosmological parameters for all the five cases and we do not detect deviations from the flat $\\Lambda$CDM paradigm. The different combinations of the four experiments used in this analysis do not show any evidence of tensions between cosmological probes. We find that the curvature is null at 1-$\\sigma$ level with errors of the order of $10^{-2}-10^{-3}$. We measure the dark energy equation of state parameter to be consistent with a cosmological constant with 13\\% uncertainty for CMB+$P_{\\mathrm{LRG}}(k)$ and 6.5\\% uncertainty for CMB+$P_{\\mathrm{LRG}}(k)$+$H_{0}$+SNIa in the flat wCDM case. If we discard the systematic errors in the SNIa, the precision increases to 4.6\\%. In the kwCDM, because of the added degree of freedom, we have a degradation of the constraints to 8.4\\%, when combining all the four samples. The constraints obtained with the CMB and large scale structure together are shifted by 1.5-2$\\sigma$ with respect to the best fit $\\Lambda$CDM because of a bias-shape degeneracy in the power spectrum that allows very large, and unphysical, values of the bias when $w_{\\mathrm{DE}}$, $\\Omega_{\\mathrm{M}}$, $\\Omega_{\\mathrm{k}}$ increase and $\\sigma_{8}$ decreases. If we assume the parametric form of equation (\\ref{eq:wa}) for the dark energy equation of state, we obtain $w_{\\mathrm{DE}}(z=0.54) = -0.97\\pm0.29$ (CMB+$P_{\\mathrm{LRG}}(k)$) and $w_{\\mathrm{DE}}(z=0.3) = -1.03\\pm0.069$ (all four experiments combined). The latter is only slightly worse than the flat wCDM result.  In the near future new and larger galaxy redshift catalogues both spectroscopic, like BOSS and HETDEX, and photometric, like Pan-STARRS and DES, will become available, together with the new measurements of the CMB anisotropies from the Planck satellite \\citep{plank_mission}. These datasets will enable us to improve the constraints presented in this work even further. Some of these experiments are explicitly designed in order to extract the maximum amount of information regarding dark energy. This would allow to exclude many classes of dark energy models or of modifications of general relativity. In order to use the full information from the power of these new large scale observations and to avoid introducing systematic effects, accurate models of the large scale distribution, scale dependent bias and redshift space distortions are necessary. In M10 we showed that the model used in this work is accurate enough to describe the power spectrum shape also for surveys with volumes larger than available nowadays. The explicit inclusion of bias and redshift space distortions in the model would allow however to use a larger range of scales than now possible, helping to improve the quality of the cosmological constrains even further."
    },
    "1107/1107.5573_arXiv.txt": {
        "abstract": "We successfully carried out the first high-altitude balloon flight of a wide-field hard X-ray coded-aperture telescope \\pe1, which was launched from the Columbia Scientific Balloon Facility at Ft. Sumner, New Mexico on October 9, 2009.  \\pe1 is the first implementation of an advanced CdZnTe (CZT) imaging detector in our ongoing program to establish the technology required for next generation wide-field hard X-ray telescopes such as the High Energy Telescope (HET) in the Energetic X-ray Imaging Survey Telescope (\\EXIST).  The CZT detector plane in \\pe1 consists of an 8 \\x 8 array of closely tiled 2 cm \\x 2 cm \\x 0.5 cm thick pixellated CZT crystals, each with 8 \\x 8 pixels, mounted on a set of readout electronics boards and covering a 256 cm\\sS{2} active area with 2.5 mm pixels.  A tungsten mask, mounted at 90 cm above the detector provides shadowgrams of X-ray sources in the 30 -- 600 keV band for imaging, allowing a fully coded field of view of 9\\Deg \\x 9\\Deg (and 19\\Deg \\x 19\\Deg for 50\\% coding fraction) with an angular resolution of 20\\arcmin. In order to reduce the background radiation, the detector is surrounded by semi-graded (Pb/Sn/Cu) passive shields on the four sides  all the way to the mask. On the back side, a 26 cm \\x 26 cm \\x 2 cm CsI(Na) active shield provides signals to tag charged particle induced events as well as $\\gtrsim$ 100 keV background photons from below. The flight duration was only about 7.5 hours due to strong winds (60 knots) at float altitude (38--39 km). Throughout the flight, the CZT detector performed excellently. The telescope observed Cyg X-1, a bright black hole binary system, for $\\sim$ 1 hour at the end of the flight. Despite a few problems with the pointing and aspect systems that caused the telescope to track about 6.4 deg off the target, the analysis of the Cyg X-1 data revealed an X-ray source at 7.2$\\sigma$ in the 30--100 keV energy band at the expected location  from the optical images taken by the onboard daytime star camera.  The success of this first flight is very encouraging for the future development of the advanced CZT imaging detectors (\\pe2, with 0.6 mm pixels), which will take advantage of the modularization architecture employed in \\pe1. ",
        "introduction": "The high energy sky is active and dynamic, full of inherently variable sources as well as extremely luminous Gamma-ray Bursts (GRBs). Wide-field hard X-ray imaging is perhaps the only effective way to capture these unpredictable, energetic celestial events. The current GRB mission, \\Swift, to study GRBS, employs a large area (0.5 m\\sS{2}) array of CdZnTe (CZT) detectors, each 4 mm \\x 4 mm \\x 2 mm as separate pixels for coded-aperture imaging. The previously proposed \\EXIST telescope would have had 4.5 m\\sS{2} with 0.6 mm pixels \\cite{Grindlay10}. A smaller (1024 cm\\sS{2}) CZT imager incorporating the same 0.6mm pixels has now been proposed for the MIRAX experiment on the {\\it Lattes} mission to be launched by Brazil \\cite{Grindlay11}. The combination of a large area and fine pixels imposes serious challenges in detector design, packaging, and operations, given the limited power and space available in space missions.  We have been pushing the technology of advanced CZT imaging detectors required for future generation of wide-field hard X-ray imagers through our ongoing, three-part, balloon-borne wide-field hard X-ray telescope experiment, \\pe.  We have successfully carried out a first high-altitude balloon flight of \\pe1, the first implementation of the program, which was launched from the NASA Columbia Scientific Balloon Facility (CSBF) at Ft. Sumner, New Mexico on October 9, 2009.  \\pe1 is a wide-field hard X-ray coded-aperture telescope, built to develop an efficient modularization and packaging architecture for a large area fine pixel CZT imaging detector (256 cm\\sS{2} with 2.5 mm pixels) and to demonstrate its performance in a near space environment. The X-ray telescope including the CZT detector performed very well during the flight.  Here we report the detailed performance of the X-ray telescope and the imaging CZT detector during the flight. Full details of the detector and telescope system, including the assembly of the coded mask, the shields, the onboard calibration source and the interface to gondola will be presented in separate papers \\cite{Allen10, Allen11}. ",
        "conclusions": "We have successfully carried out the first high altitude (40 km) balloon flight of the wide-field hard X-ray telescope, \\pe1, employing the first generation of an advanced CZT imager. The CZT imager in \\pe1 is,  to our knowledge, the largest pixellated close-tiled CZT array at the moment (256 cm\\sS{2} with 2.5 mm pixels).  The X-ray telescope performed excellently throughout the 7.5 hr flight.  Despite a few problems in the new aspect and pointing system, we were able to detect Cyg X-1 at 7.2$\\sigma$ from about an hour observation, $\\sim$ 50 min of which was with relatively poor aspect and so with smearing in the elevation axis direction.  Encouraged by this success, we have begun the next phase of the program, \\pe2, which will have 4\\x finer detector spatial resolution (pixel pitch: 0.6 mm) and also both lower energy threshold ($\\sim$ 5 -- 8 keV) and better energy resolution ($\\sim$ 2 keV). To accomplish this, a new ASIC, the Nu-ASIC, developed for \\NuSTAR\\footnote{\\url{http://www.nustar.caltech.edu/}} will be used in the 8\\x 8 DCU close-tiled CZT detectors of \\pe2. The Nu-ASIC is developed in the same ASIC family line as the RadNET ASIC, so they share a similar basic operational architecture, which simplifies the transition. We will combine the Nu-ASIC front-end readout electronics with a new packaging and modularization architecture derived from the \\pe1 system. For more details of the \\pe2 detector development, see also \\cite{Hong10}."
    },
    "1107/1107.4739_arXiv.txt": {
        "abstract": "We study the hybrid inflationary potential in a regime where the defect field is light, and more than 60 e-folds of accelerated expansion occur after the symmetry breaking transition. Using analytic and numerical techniques, we then identify parameter values within this regime for which the statistics of the primordial curvature perturbation are significantly non-Gaussian.  Focusing on this range of parameters, we provide a specific example which leads to an observationally consistent power spectrum, and a level of non-Gaussianity within current WMAP bounds and in reach of the Planck satellite. An interesting feature of this example is that the initial conditions at horizon crossing appear quite natural. ",
        "introduction": "An inflationary epoch is the most promising model for the origin of large-scale structure in the universe \\cite{inflation} (for reviews see \\cite{reviews}), naturally accounting for current observations which indicate that the universe's structure originated from near-scale invariant and almost Gaussian fluctuations \\cite{observations,wmap}. The small deviations from scale-invariance and Gaussianity are key probes of inflation, allowing competing models to be distinguished from one another by present and future observations.  In canonical, single field inflation, the curvature perturbation produced at horizon crossing is extremely close to Gaussian, and is conserved through the entirety of the subsequent evolution \\cite{Maldacena:2002vr,Bardeen,Lyth,zetaConv}. On the other hand,  when more than one field is light at horizon crossing, isocurvature modes are produced, which later can source an additional contribution to the curvature perturbation. This can cause the curvature perturbation to become sufficiently non-Gaussian for this signal to be detected by future observations, or the model constrained by present ones. The current WMAP constraints on the $\\fnl$ parameter, which measures the deviation of the three--point correlation function from zero, is $-10<\\fnl <74$ at $95\\%$ confidence level, but $\\fnl \\approx 32\\pm 21$ at the $68\\%$ level~\\cite{wmap}. Though premature, it is tempting to consider these numbers as a `tentative hint' of a non-Gaussian signal. The recently launched PLANCK satellite, will help resolve the issue. In the absence of a detection it is expected to give the constraint $|\\fnl|<5$.  With these issues in mind, considerable effort has been invested in developing methods to calculate how non-Gaussianity evolves on super-horizon scales during inflation, \\cite{Lyth:2005fi,Vernizzi:2006ve, otherMethods}, into identifying the features an inflationary potential must possess and the initial conditions required, to lead to a `large' non-Gaussianity ($\\fnl>1$) \\cite{largeNG1,largeNG2,Kim:2010ud,Elliston:2011dr}. Many scenarios have also been studied in which the generation of non-Gaussianity occurs due to post inflationary effects, such as in curvaton \\cite{curv}, preheating~\\cite{preh}, modulated (p)reheating \\cite{mod}, and models with an inhomogenious end of inflation \\cite{inHom1,inhom2} (see Ref.~\\cite{Alabidi:2010ba} for connections between these three scenarios).  Of particular interest are effects which occur for inflationary potentials which are well motivated from a particle physics perspective. The hybrid potential is one such example, and moreover is one of the most widely studied potentials in the literature. First introduced as an inflationary model by Linde \\cite{Linde:1993cn}, it is a two-field potential of the form \\be V(\\phi,\\chi)=\\frac{1}{2}m^2\\phi^2+\\frac{1}{2}g^2\\phi^2\\chi^2+\\frac{\\lambda}{4}(\\chi^2-v^2)^2\\,. \\label{eq:hybridPot} \\ee As a quantum field theory, this is attractive because it is renormalisable and does not require unnatural cancellation of radiative corrections. In fact, potentials of this type are common in supersymmetric models \\cite{SUSY}.   In this paper we will investigate the hybrid potential (\\ref{eq:hybridPot}) using a different choice of parameters from the original proposal. We find that this leads to a viable cosmological model, with perturbations that are compatible with current observations and level of non-Gaussianity that would be observable with the Planck satellite. ",
        "conclusions": "In this paper we have demonstrated that inflation sourced by the hybrid potential Eq.~(\\ref{eq:hybridPot}) can give rise to an observationally significant non-Gaussianity while satisfying WMAP constraints on the power spectrum. The conditions required for this are summarised in Eq.~(\\ref{eq:ourParameters}). This is an example of the hilltop mechanism for producing non-Gaussianity first identified for axion models in Ref.~\\cite{Kim:2010ud}.  In contrast with original hybrid inflation, where the waterfall field is heavy and cannot contribute significantly to the curvature perturbation on large scales \\cite{heavy}, we take the waterfall field to be light. Our scenario also leads to more than $60$ e-folds of accelerated expansion after the waterfall transition. This avoids issues of primordial black holes and domain walls. While earlier work has considered the hybrid model with a light waterfall field and the resulting non-Gaussianity \\cite{barnaby, Abolhasani:2011yp} (and hybrid inflation with two light inflaton fields \\cite{largeNG1,inhom2}), we believe that our work is the only study of non-Gaussianity in the case in which more than $60$ e-folds of evolution occurs after the hybrid transition.  Numerically we have only exhaustively studied the range of parameters discussed in Section~\\ref{sec:numerical}, but we have also probed several other parameter choices which lead to $60$ e-folds after the hybrid transition, including a number which are outside of the regime summarised by Eq.~(\\ref{eq:ourParameters}). Inside this regime analytic expectations are confirmed remarkably well by our simulations, which suggests that we should be able to trust them even outside it as long as the assumptions thay are based on are satisfied. Outside of this regime our non-exhaustive probing has not found any parameter choices which lead to a significant non-Gaussianity. It would be interesting, however, to probe an extended parameter range more carefully in order to determine whether the conditions Eq.~(\\ref{eq:ourParameters}) are the only ones which lead to a signal.  Finally we note that as discussed in Section~\\ref{sec:ics} an interesting feature of the hybrid model is that it offers an explanation for the initial conditions at horizon crossing which lead to a large non-Gaussianity. While parameter choices are necessarily fine tuned, therefore, the initial conditions need not be, at least for some subset of the parameters which can produce a large non-Gaussianity. This is in contrast to other two field models which produce a large non-Gaussianity due to inflationary dynamics \\cite{largeNG1, largeNG2,Elliston:2011dr}, which require both carefully choosen model parameters, and offer no explanation for the origin of the finely tuned horizon crossing conditions."
    },
    "1107/1107.1093_arXiv.txt": {
        "abstract": "{The blazar \\object{3C 454.3} is one of the most active sources from the radio to the $\\gamma$-ray frequencies observed in the past few years.} {We present multiwavelength observations of this source from April 2008 to March 2010. The radio to optical data are mostly from the GASP-WEBT, UV and X-ray data from Swift, and $\\gamma$-ray data from the AGILE and Fermi satellites. The aim is to understand the connection among emissions at different frequencies and to derive information on the emitting jet.} {Light curves in 18 bands were carefully assembled to study flux variability correlations. We improved the calibration of optical--UV data from the UVOT and OM instruments and estimated the Ly$\\alpha$ flux to disentangle the contributions from different components in this spectral region.} {The observations reveal prominent variability above 8 GHz. In the optical--UV band, the variability amplitude decreases with increasing frequency due to a steadier radiation from both a broad line region and an accretion disc. The optical flux reaches nearly the same levels in the 2008--2009 and 2009--2010 observing seasons; the mm one shows similar behaviour, whereas the $\\gamma$ and X-ray flux levels rise in the second period. Two prominent $\\gamma$-ray flares in mid 2008 and late 2009 show a double-peaked structure, with a variable $\\gamma$/optical flux ratio. The X-ray flux variations seem to follow the $\\gamma$-ray and optical ones by about 0.5 and 1 d, respectively. } {We interpret the multifrequency behaviour in terms of an inhomogeneous curved jet, where synchrotron radiation of increasing wavelength is produced in progressively outer and wider jet regions, which can change their orientation in time. In particular, we assume that the long-term variability is due to this geometrical effect. By combining the optical and mm light curves to fit the $\\gamma$ and X-ray ones, we find that the $\\gamma$ (X-ray) emission may be explained by inverse-Comptonisation of synchrotron optical (IR) photons by their parent relativistic electrons (SSC process). A slight, variable misalignment between the synchrotron and Comptonisation zones would explain the increased $\\gamma$ and X-ray flux levels in 2009--2010, as well as the change in the $\\gamma$/optical flux ratio during the outbursts peaks. The time delays of the X-ray flux changes after the $\\gamma$, and optical ones are consistent with the proposed scenario. } ",
        "introduction": "A relativistic jet pointing at a small angle to the line of sight is most likely responsible for the extreme properties of the active galactic nuclei known as ``blazars\", i.e.\\ BL Lac objects and flat-spectrum radio quasars (FSRQ). Indeed, the alignment would cause Doppler enhancement of the emission and contraction of its variability time scales. This peculiar orientation would also explain the apparent superluminal motion of radio knots in the jet. Owing to the jet emission beaming, we observe intensified synchrotron radiation from the radio to the UV--X-ray band and inverse-Compton radiation at higher energies, up to the TeV domain. Sometimes, in low brightness states, other contributions are detected, likely because of  unbeamed radiation from the blazar nucleus, i.e.\\ the accretion disc and broad line region (BLR). This occurs more often in FSRQ than in BL Lac objects, and indeed the classical distinction between the two classes is based on the equivalent width of their broad emission lines (that has to be greater than 5 \\AA\\ in the rest frame for FSRQ). One of the main issues concerns the origin of the seed photons that are Comptonised to X- and $\\gamma$-ray energies: relativistic electrons certainly upscatter the soft photons they have previously produced by synchrotron emission (synchrotron-self-Compton, or SSC, process), but possibly also photons coming from outside the jet (external Compton, or EC, process), in particular from the accretion disc, the BLR, or an obscuring torus \\citep[see e.g.][and references therein]{der09}.  The quasar-type blazar \\object{3C 454.3} has received particular attention by the international astronomical community since it underwent a major outburst in 2005 \\citep{vil06,gio06,pia06,fuh06}. Indeed, after many decades of intense radio, but only mild optical activity, the source began brightening in the optical band in 2001, until in May 2005 it reached the maximum optical level ever observed, $R=12.0$. The outburst was simultaneously detected also in X-rays, while the millimetric radio flux peaked about one month after, and the 15--43 GHz flux much later \\citep{vil07,rai08c}. According to \\citet{vil06,vil07} the observed change in behaviour starting from 2001 was a geometric effect, i.e.\\ the motion of a curved jet producing variations in the viewing angle of the different emitting regions.  In the following 2006--2007 observing season the source remained in a faint state, and new features appeared in its spectral energy distribution (SED). Through the analysis of low-energy data from the Whole Earth Blazar Telescope\\footnote{\\tt http://www.oato.inaf.it/blazars/webt/} (WEBT) and high-energy data from the XMM-Newton satellite, \\citet{rai07b} were able to recognise both a little blue bump in the optical, possibly due to emission lines from the BLR, and a UV excess, suggesting a big blue bump likely due to thermal radiation from the accretion disc. Moreover, they argue that the X-ray spectrum may be concave, with spectral softening at lower energies becoming more evident when the source is fainter, maybe revealing the high-frequency tail of the big blue bump. Finally, they ascribed the flux excess in the $J$ band to a prominent broad H$\\alpha$ emission line. Both the UV excess and X-ray spectral curvature were also inferred from further XMM-Newton observations, at the beginning of the following observing season, while the contribution of the H$\\alpha$ line was confirmed by near-IR spectroscopy at Campo Imperatore  \\citep{rai08c}.  Then, from July 2007 the source began a new activity phase, during which it was detected several times in $\\gamma$-rays by the AGILE satellite\\footnote{\\tt http://agile.iasf-roma.inaf.it/}, and was monitored from the optical to the radio bands by the WEBT and its GLAST-AGILE Support Program \\citep[GASP,][]{vil08}. The results of these observations were published in \\citet{ver08}, \\citet{rai08b,rai08c}, \\citet{ver09}, \\citet{don09b}, \\citet{vil09b}, and \\citet{ver10}. Rotations of the optical polarisation vector both clockwise and counterclockwise were detected by \\citet{sas10}. In contrast, observations at very high energies ($\\rm E > 100 \\, GeV$) by the Cherenkov telescope MAGIC resulted in only upper limits \\citep{and09}.  From the theoretical point of view, \\citet{ghi07} compare the source SED of July 2007 with that of May 2005 and infer that the dissipation site in relativistic jets changes, with more compact emitting regions, smaller bulk Lorentz factors, and greater magnetic fields characterising locations closer to the black hole. However, in the \\citet{ghi07} view, the dissipation always occurs in a subparsec zone, where the seed photons for Compton emission are provided by the BLR. In contrast, \\citet{sik08} claim that the dissipation region is located much farther, in the millimetric photosphere, at some parsecs from the central engine, so the Compton emission is due to scattering of infrared photons emitted by the hot dust.  In the meanwhile, the Fermi satellite\\footnote{\\tt http://fermi.gsfc.nasa.gov/} was launched, providing an unprecedented monitoring of 3C 454.3 at $\\gamma$-ray energies \\citep{bon09,abd09_3c454}. The data from AGILE and Fermi were separately used to study the correlations between flux variations in optical and $\\gamma$-rays. All results agreed that the time lag of $\\gamma$-ray flux variations after the optical ones is less than one day \\citep{ver09,don09b,bon09,ver10}. \\citet{fin10} discuss the spectral break at $\\sim 2.4 \\rm \\, GeV$ shown by Fermi data in August 2008 suggesting that it results from the combination of Compton-scattered disc and BLR radiations.  An analysis of the radio--optical data in 2005--2007, including space observations by Spitzer in the infrared, revealed two synchrotron peaks, the primary at IR and the secondary at sub-mm wavelengths, likely arising from distinct regions in the jet \\citep{ogl10}.  A multiwavelength study of the flaring behaviour of 3C 454.3 during 2005--2008 was performed by \\citet{jor10}. These authors suggest that the emergence of a superluminal knot from the core produces a series of optical and high-energy outbursts, and that the millimetre-wave core lies at the end of the jet acceleration and collimation zone.  Suzaku observations of 3C 454.3 in November 2008 analysed by \\citet{abd10_suzaku} confirmed the earlier suggestions by \\citet{rai07b,rai08c} that there seems to be a soft X-ray excess, which  becomes more important when the source gets fainter. They interpret it as either a contribution of the high-energy tail of the synchrotron component or bulk-Compton radiation produced by Comptonisation of UV photons from the accretion disc by cold electrons in the innermost parts of the relativistic jet.  More recently, \\citet[][see also \\citealt{str10}]{pac10} have reported on a $\\gamma$-ray flare of 3C 454.3 detected by AGILE in December 2009, when the peak flux reached $2.0 \\times 10^{-5} \\rm \\, photons \\, cm^{-2} \\, s^{-1}$. The simultaneous flux increase at UV and optical frequencies was by far less dramatic. The authors found that the source behaviour during the flare cannot be accounted for by a simple one-zone model, but an additional particle component is needed. In contrast, \\citet{bon11} explain the source behaviour in the same period by means of a one-zone model, where the magnetic field is slightly weaker when the overall jet luminosity is higher. A new intense $\\gamma$-ray flare was detected by Fermi in April 2010 \\citep[$\\approx 1.6 \\times 10^{-5} \\rm \\, photons \\, cm^{-2} \\, s^{-1}$,][]{ack10}, when fast variability on a time scale of a few hours was revealed \\citep{fos10}, supporting earlier results by \\citet{tav10}, who claim that the dissipation region must be close to the black hole. The December 2009 and April 2010 $\\gamma$-ray outbursts detected by Fermi were further analysed by \\citet{ack10}. These authors in particular discuss the observed spectral break at a few GeV, which they ascribe to a break in the underlying electron spectrum. This feature was also observed during the unprecedented $\\gamma$-ray outburst of November 2010 \\citep{abd11}, when the $\\gamma$ flux rose to $(6.6 \\pm 0.2) \\times 10^{-5} \\rm \\, photons \\, cm^{-2} \\, s^{-1}$ showing clear spectral changes.  In this paper we present the radio-to-optical monitoring results obtained by the GASP-WEBT in 2008--2010, together with UV and X-ray data acquired by Swift and $\\gamma$-ray observations by the AGILE and Fermi satellites. We analyse the flux correlations at different frequencies to derive information on the jet physics and structure. As in several previous WEBT-GASP papers, we focus on an interpretation of the observed light curves in terms of variations in the orientation and curvature of an inhomogeneous emitting jet. ",
        "conclusions": "In this paper we have presented multifrequency observations of 3C 454.3 from April 2008 to March 2010. Radio-to-optical data are mainly from the GASP-WEBT, and UV and X-ray data are from Swift and $\\gamma$-ray data from AGILE and Fermi. In these two observing seasons, the source showed prominent outbursts at frequencies greater than 8 GHz, with rapid flares superposed in the optical-to-$\\gamma$ light curves.  We have reanalysed the question of the nature of the optical--UV emission, deriving more accurate calibrations for both the UVOT instrument onboard Swift and the OM detector onboard XMM-Newton, and estimating the contribution of the emission lines (mainly Ly$\\alpha$) in the UV. We confirmed that the observed decrease of the variability amplitude with increasing frequency from the near-IR to the UV band is due to the contribution of radiation from both the BLR and the accretion disc. The long-term flux variations from the optical to the $\\gamma$-ray frequencies appear roughly correlated. The mm flux remained high during the whole optical flaring period.  By comparing the multiwavelength behaviour in the 2008--2009 and 2009--2010 observing seasons, we noticed that the optical flux reached nearly the same levels (i.e.\\ minimum and maximum brightness) in both seasons. The same thing occurred in the mm band, whereas the $\\gamma$-ray and X-ray flux levels increased in the second period.  We interpreted the observations in terms of an inhomogeneous jet model, where synchrotron radiation of increasing wavelength is produced in progressively larger regions farther away from the emitting jet apex. The jet is supposed to be a curved (possibly in a helical shape) and flexible (maybe rotating) structure, where different emitting zones have different alignments with the line of sight, and their viewing angle can change in time. A similar model has been adopted to explain the multiwavelength behaviour of Mkn 501 \\citep{vil99}, S4 0954+65 \\citep{rai99}, BL Lacertae \\citep{vil02,vil04a,vil09a,rai09,rai10,lar10}, AO 0235+164 \\citep{ost04}, along with 3C 454.3 \\citep{vil06,vil07,vil09b}.  We assumed that the long-term flux density variations, with a tentative time scale of one week in the optical and of one month in the mm band, are due to changes of the viewing angles of the corresponding emitting regions, while variability on shorter time scales can be ascribed either to intrinsic flaring activity or to a finer geometrical structure (as suggested by \\citealt{lar10} for BL Lacertae). Then, under the further assumption that the $\\gamma$ and X-ray emissions are produced by inverse-Compton scattering of synchrotron photons and that the behaviour of the IR light curve is intermediate between the optical and mm ones, we combined the optical and mm fluxes to fit the $\\gamma$ and X-ray light curves. In this scenario the multiwavelength observations presented in this paper suggest that the $\\gamma$-ray radiation comes from inverse-Comptonisation of ``optical\" (actually UV to-near-IR) synchrotron photons produced in the inner emitting jet region by their parent relativistic electrons (SSC process). However, depending on the mean free path before upscattering, the bulk of the $\\gamma$-ray radiation is produced in a region slightly downstream of the one responsible for the bulk of the optical emission. The jet curvature then implies that the $\\gamma$ and optical regions are slightly misaligned. A reduction of this misalignment, allowing more optical seed photons to enter the Comptonisation zone, could thus be at the origin of the $\\gamma$-ray flux increase in 2009--2010. Moreover, the variation in the $\\gamma$/optical flux ratio during the outburst peaks of mid 2008 and late 2009 can be explained by the fact that the $\\gamma$ zone acquires the minimum viewing angle, hence the maximum Doppler boosting, some days before the optical zone.  The X-ray light curve is not as well sampled as the optical and $\\gamma$-ray ones, which makes a detailed analysis more difficult. However, the above-mentioned fit suggests that the X-ray radiation comes from a downstream region, where relativistic electrons produce and upscatter synchrotron IR photons (SSC process). The rise in the X-ray flux levels in the 2009--2010 season would then be explained similarly to the case of the $\\gamma$ flux, with a better alignment between the zones responsible for the bulk of the synchrotron and inverse-Compton emissions. The mean delay of about 1 d of the X-ray flux variations after the optical ones and of about 0.5 d after the $\\gamma$-ray ones, found by cross-correlation analysis, are consistent with this picture."
    },
    "1107/1107.4480_arXiv.txt": {
        "abstract": "Using data from the Extreme-ultraviolet Imaging Spectrometer aboard Hinode, we have studied the coronal plasma in the core of two active regions. Concentrating on the area between opposite polarity moss, we found emission measure distributions having an approximate power-law form EM$\\propto T^{2.4}$ from $\\log\\,T = 5.5$ up to a peak at $\\log\\,T = 6.55$.  We show that the observations compare very favourably with a simple model of nanoflare-heated loop strands. They also appear to be consistent with more sophisticated nanoflare models.  However, in the absence of additional constraints, steady heating is also a viable explanation. ",
        "introduction": "It has proved very challenging to obtain definitive answers to the very stubborn problem of how the solar corona is heated. The heating of warm ($\\sim 1$ MK) coronal loops is one exception. There is now widespread agreement, though not complete consensus, that these are multi-stranded structures heated impulsively by storms of nanoflares (e.g., Klimchuk 2006, 2009 and references cited therein). Warm loops account for only a fraction of the coronal plasma, however. The corona also contains hot ($> 2$ MK) loops and diffuse emission at all temperatures.  Much recent attention has focused on what is often called the ''hot core\" of active regions. This is the subject of our study reported here.  A fundamental question concerns the temporal nature of coronal heating. This is often discussed in terms of a dichotomy between steady heating and nanoflares. What really matters for determining the properties of the coronal plasma is the heating on individual magnetic flux strands (field lines).  Few, if any, serious coronal heating theories predict that the energy release is steady in this regard \\citep{klimchuk_2006}; rather, the heating occurs in short-lived bursts. We refer to these small-scale impulsive bursts as nanoflares. No particular physical mechanism is implied by the term. It could be magnetic reconnection in tangled magnetic fields as envisioned by \\cite{parker_1983} and probably involving the secondary instability of current sheets \\citep{dahlburg_2005}; it could be wave dissipation in drifting resonance layers \\citep{ofman_1998}; or it could be another mechanism altogether.  An important parameter is the time delay between successive nanoflares on a given strand. By delay we mean the time interval between the end of one nanoflare and the start of the next.  If the delay is much shorter than a cooling time (tens to hundreds of seconds) then there is minimal cooling between events and the heating is effectively steady. We can approximate such a situation with perfectly steady heating. It is becoming more common for the term ''steady heating\" to imply nanoflares that repeat with a high enough frequency that the temperature hovers around one value and for ''nanoflare heating\" to imply nanoflares that repeat with a low enough frequency that the plasma cools substantially. In addition, ''nanoflare heating\" usually implies heating that takes place in the corona. Impulsive energy release in the chromosphere that is associated with spicules \\citep{depontieu_2011} is a different phenomenon, though it might have similar physical attributes.  Attempts to determine whether active region cores have steady or nanoflare heating have been inconclusive.  Some studies point out that the intensities, electron densities, Doppler shifts, and nonthermal broadening of observed emissions are often quite steady in the sense that fluctuation amplitudes are small \\citep{antiochos,tripathi_moss_2010,brooks_warren}. A reasonable conclusion is that the heating is steady.  However, this is only certain if the cross-field spatial scale of the heating is comparable to or larger than the resolution of the observations (typically about 1000 km). There is good reason to believe that the coronal magnetic field is structured on a much smaller scale \\citep{klimchuk_2006} and so it seems likely that the heating is also structured on a much smaller scale. If so, direct evidence of nanoflares would be washed out, since the emission from many different strands would be averaged together. Such averaging occurs both across the observational pixel and along the optically thin line of sight. Hence, steady intensities, densities, Doppler shifts, and nonthermal broadening are consistent with both steady heating and nanoflare heating.  We note also that subtle variability in seemingly steady intensities can provide good evidence for nanoflares (Terzo et al. 2011; Viall \\& Klimchuk 2011b).  Faced with the likelihood of unresolved structures, we must look to other ways of distinguishing between steady and impulsive heating. Investigating the distribution of plasma with temperature, as quantified in the the emission measure distribution, EM($T$), is one promising approach.  There have been many determinations of EM($T$) through the years, but they have tended to average over large areas. Results reported for active regions generally obey the power law EM($T$) $\\propto T^b$ up to a peak near 3 MK, where the index $b$ ranges between 1 and 3 (e.g., \\citealt{dere_1982,dere_1993,brosius_1996}). When plotted on a log-log scale, $b$ is the slope of a straight line.  Some authors use the differential emission measure, DEM($T$), which is related to the emission measure by EM($T$) = $T \\times$ DEM($T$).  The slope of DEM($T$) is smaller by an amount 1.0.  Observations that average over large areas tend to include both coronal and transition region emissions.  Transition region here refers not to a particular temperature range, but rather to the thin region of steep temperature gradient at the base of the corona.  As a rule of thumb, the temperature of the top of the transition region is about 60\\% of the maximum temperature in the strand \\citep{ebtel,bradshaw_2010}.  The transition region can therefore reach very high temperatures depending on how hot the strand is. Million degree moss seen in the 171 channels of SOHO/EIT, TRACE, and SDO/AIA is just the transition region footpoints of hot core strands. We use this later to determine the width of the core in the active region we investigate here.  Moss can tell us a great deal about the core plasma. For example, in a recent paper we showed that the EM distribution of moss is better explained if the core is heated impulsively than in a steady fashion \\citep{tripathi_moss_2010_apj}. Our results are compatible with steady heating if the cross sectional area has the proper temperature dependence in the transition region ''throat\" of a rapidly expanding strand, but we argued that this temperature dependence is not likely to be maintained in the presence of spatial and temporal pressure nonuniformities across the moss.  Two recent studies have explicitly avoided moss and concentrated on the coronal emission in the cores of active regions. \\citet{warren_2011} studied a small ($10\\times15$ arcsec$^2$) inter-moss region between opposite magnetic polarities and found an emission measure distribution that can be approximated by EM $\\propto T^{3.26}$ in the range $6.0 \\le \\log\\,T \\le 6.6$. The EM is 400 times smaller at $\\log\\,T = 5.8$ than it is at $\\log\\,T = 6.6$. \\citet{winebarger_2011} averaged over a $5\\times25$ arcsec$^2$ inter-moss area from another active region and found a similar power-law slope ($EM \\propto T^{3.2}$) over the range $6.0 \\le \\log\\,T \\le 6.5$. These distributions are considerably steeper than what has been published in the past. Is this because the previous studies averaged over a mixture of different features (core, moss, non-core loops, etc.), or is it because the inter-moss regions of the two new studies are, by chance, not typical of core plasma in general?  To help answer this question, we have analyzed four inter-moss regions within active region AR 10961 that was observed by multiple spacecraft and one small inter-moss region in active region AR 10980. In the next sections, we describe the observations and our analysis of the data. We derive emission measure distributions for the inter-moss regions, and by estimating the plasma in the foreground of the core, we determine the EM distribution of the core plasma itself. We then discuss the results in the context of related measurements and the predictions of steady and nanoflare heating. ",
        "conclusions": "\\label{discussion}  The ultimate goal is to understand what these observations are telling us about coronal heating.  Is the heating effectively steady or does it take the form of low-frequency nanoflares with substantial cooling between events?  A third possibility is that plasma is heated in the chromosphere, not the corona, and ejected upward as a spicule  \\citep{depontieu_2011}, although see \\citet{klimchuk_2011}.  We concentrate here on the first two possibilities, since the spicule explanation is newly developing, and many details are yet to be worked out.  If an observed inter-moss region is small enough that only short segments of strands (loops) are observed, then any shape of emission measure distribution is consistent with steady heating. Since the strands are not evolving, the only requirement is to have the right number of strands at each temperature.  We must then rely on other information to determine whether this distribution of strands is feasible. Fortunately, strands that are at or near static equilibrium have well known and highly restrictive physical constraints.  If the heating is not too asymmetric \\citep{winebarger_2002,patsourakos_etal_2004} and not too highly concentrated near the footpoints \\citep{klimchuk_2010}, then the coronal temperature is uniquely related to the pressure and length \\citep{rosner_1978,craig_1978}. \\citet{winebarger_2011} made use of this in an impressive modeling effort that was a key part of their study. From density sensitive line ratio measurements at the moss footpoints and a potential magnetic field extrapolation, they determined the pressures and lengths of the core strands that pass through the LOS in the observed inter-moss region.  They then determined the temperatures of the strands by assuming static equilibrium and eventually obtained an EM distribution. The distribution they found in this way is considerably more narrow than what they observed. It matches the observations reasonably well in the range $6.3 \\le \\log\\,T \\le 6.7$, but it drops abruptly to zero at the extremes, unlike the observations. \\cite{winebarger_2011} suggest that the excess emission observed at warm temperatures is due to foreground plasma above the core. They subtracted an estimate of the foreground, but there is still much more warm EM in the inter-moss region than the model predicts. It is possible that the foreground was underestimated; on the other hand, no account was made for gravitational stratification, which would tend to produce an overestimate.  As mentioned earlier, treating the foreground is a difficult business.  What about low-frequency nanoflares?  In a simple attempt to model the core arcade of our active region, we performed a nanoflare simulation using the EBTEL hydrodynamics code \\citep{ebtel}.  We considered one representative strand (loop) with a half length of $2.4 \\times 10^9$ cm.  For a semi-circular shape, this corresponds to a radius of $1.5 \\times 10^9$ cm, or half the estimated radius of the arcade.  Thus, the strand has something of an average length for all the strands in the arcade.  The nanoflares have a triangular heating profile with a duration of 500 s and amplitude of 0.04 erg cm$^{-3}$ s$^{-1}$. There is also a constant low-level background heating of $10^{-6}$ erg cm$^{-3}$ s$^{-1}$. The nanoflares repeat every 8000 s.  In order to obtain a predicted EM distribution for an inter-moss observation, we assume that the time-averaged properties of the coronal part of the model strand apply throughout a vertical column of length $3 \\times 10^9$ cm.  This is equal to the height of the arcade. The time-averaged energy flux needed to maintain the column is $3.75 \\times 10^6$ erg cm$^{-2}$ s$^{-1}$, which is a very reasonable value for active regions \\citep{withbroe_1977}.  The solid curve in Fig.~\\ref{model} shows the EM distribution from our nanoflare model. The agreement with the observations is excellent. Nearly all of the data points deviate less than the $\\Delta \\log EM \\approx \\pm 0.2$ uncertainty. Using this EM distribution we have also predicted intensities for different spectral lines shown in Table~\\ref{compare}. The predicted intensities are within 20-30\\% to the observed for most of the lines.  It is far too premature, however, to conclude that core is necessarily heated by nanoflares. Before any such conclusions can be drawn, we must construct a more realistic and complete model, and we must consider additional observational constraints, as was done by \\cite{winebarger_2011} for steady heating.  Such a modeling effort is already underway. We also caution that the original EBTEL code is most accurate during the heating and conductive cooling phases of the strand evolution, and least accurate during the radiation/enthalpy cooling phase. It tends to underestimate the emission measure of the cooler plasma and therefore to overestimate the slope of the EM distribution. We are in the process of making improvements to the code \\citep{cargill_2011}.  \\citet{mulu-moore_2011} used the more sophisticated NRLFTM hydrodynamics code to model nanoflare heating and found EM slopes ranging from 2.0 to 2.3, depending on the loop length and nanoflare energy (assuming coronal abundances). A composite model made by combining the individual cases has a slope of 2.0. \\citet{warren_2011} used the same code to predict a hot-to-warm emission measure ratio similar to what we observe, but a factor of 10 smaller than what they observe.  Finally, Bradshaw \\& Klimchuk (2011, private communication) performed nanoflare simulations that include the effects of non-equilibrium ionization and found slopes consistent with those of \\citet{mulu-moore_2011}.  They also found that the slope increases with nanoflare duration and energy, as did \\citet{ebtel}.  As a final caveat, we note that the hottest plasma in a nanoflare heated strand could deviate greatly from ionization equilibrium. The intensities of emission lines such as the Fe XVII line in our study may then be suppressed compared to what ionization equilibrium would imply \\citep{reale_2008, bradshaw_2011}. If that is case here, then the actual emission measure at the hot extreme of our observed distribution is greater than indicated in Fig.~\\ref{model}. More intense nanoflares would be needed to reproduce the observations.  In summary, we have studied the inter-moss cores of active regions AR 10961 and AR 10980 and found  emission measure distributions with a power law slope of approximately 2.4.  This is comparable to, but on the steep end of what has been reported previously for active regions as a whole.  However, it is considerably less steep than what has been found recently in two other inter-moss regions \\citep[][]{warren_2011, winebarger_2011}. This raises an important question as to what is typical. Future investigations of emission measure distributions in the cores of many active regions are needed before we have a clear answer. Our observations are in good agreement with the predictions of a simple nanoflare model. However, in the absence of additional observational constraints such as pressure, plasma flows and non-thermal velocities, the observations could equally well be explained by steady heating. Future studies should include these additional constraints if possible."
    },
    "1107/1107.0576_arXiv.txt": {
        "abstract": "We used deep wide-field observations obtained with the Canada-France-Hawaii Telescope to study the blue straggler star (BSS) population in the innermost five arcminutes of the remote Galactic globular cluster Palomar 14. The BSS radial distribution is found to be consistent with that of the normal cluster stars, showing no evidence of central segregation.  Palomar 14 is the third system in the Galaxy (in addition to $\\omega$ Centauri and NGC 2419) showing a population of BSS not centrally segregated. This is the most direct evidence that in Palomar 14 two-body relaxation has not fully established energy equipartition yet, even in the central regions (in agreement with the estimated half-mass relaxation time, which is significantly larger than the cluster age). These observational facts have important implications for the interpretation of the shape of the mass function and the existence of the tidal tails recently discovered in this cluster. ",
        "introduction": "\\label{sec_intro}  Blue straggler stars (BSS) are hydrogen-burning stars located bluer and brighter than the main sequence (MS) turn-off (TO) point in the optical color-magnitude diagram (CMD) of star clusters.  There is a general consensus in considering BSS as the most massive (with $M_{\\rm BSS}=1-1.4M_{\\odot}$; see Shara et al. 1997, Ferraro et al 2006a) luminous objects in the CMD of a globular cluster (GC). Two physical mechanisms (both affecting and affected by the dynamical processes occurring in the cluster) have been proposed for their formation: direct stellar collisions and mass transfer activity in binary systems \\citep{colbss, mtbss1}.  The former is expected to be particularly important in high-density environments, while it should be less efficient in loose GCs and in the cluster external regions, with respect to the undisturbed evolution of primordial binaries.  On the other hand, the discovery \\citep{fe09} of two distinct and parallel BSS sequences in M30, possibly populated by objects generated by the two formation processes during the core collapse phase, further supports the idea that BSS of different origins can coexist within the same stellar system.  Observational proof of the connection between the binary and the BSS populations in low-density environments is testified by the correlation between the number fraction of these two species recently found in the core of 13 loose GCs \\citep{so08}. In addition, being more massive than the average, BSS tend to sink toward the bottom of the cluster potential well, under the action of energy equipartition. In most of the surveyed GCs ($\\sim20$ to date) BSS appear to be strongly concentrated in the core (see, e.g., Figure 2 in Ferraro et al. 2003 and Figure 6 in Ferraro \\& Lanzoni 2009), with their fraction, measured with respect to ordinary cluster stars, decreasing at increasing distance from the center and showing an upturn in the external regions \\citep[see also][]{fe93, fe04,dal09}. This feature has been interpreted by means of numerical simulations~\\citep[][]{ma04,ma06,la07a,la07b} showing that the central peak is due both to BSS formed in place because of stellar collisions, and to mass transfer BSS sunk to the centre under the effect of mass segregation.  In contrast, the rising branch in the cluster outskirts is due to BSS generated by the unperturbed evolution of primordial binaries, which are preferentially orbiting in regions where the dynamical friction timescale is longer than the cluster age. The only two exceptions currently known are $\\omega$Centauri \\citep{fe06} and NGC 2419 \\citep{da08}, where BSS show the same radial distribution as that of the other cluster stars. This fact indicates that these two GCs have not yet reached a status of energy equipartition, and their BSS result from the evolution of binary systems whose radial distribution has not been altered by the the process of mass segregation.  Based on these results, here we use the BSS radial distribution to investigate the dynamical state of the remote GC Palomar 14 (hereafter Pal14), located in the outer Halo of the Milky Way, at a distance of $\\sim71$ kpc \\citep[][hereafter S11]{so11}.  Pal14 has been indicated as one of the best candidates to test alternative theories of gravity \\citep{bau05, solnip10} and recent analyses have shown a peculiar structure and kinematics of this cluster. Adopting a sample of 17 stars \\citet{jo09} measured a very small velocity dispersion ($\\sigma_{v}=0.38\\pm0.12$ km/s). \\citet{ku10} claimed that such a small value is not compatible with the presence of the fraction of binaries ($>20\\%$) expected in a loose GC. They argued that either this cluster hosts a low fraction of binaries, or it constitutes a ``deep-freeze'' with an unusually low velocity dispersion.  By measuring the cluster mass function between 0.53 and 0.78$M_\\odot$, \\citet{jo09} found a significantly flatter slope than the canonical value and concluded that either Pal14 formed with only few low-mass stars, or it is mass segregated and lost most of its low-mass stars through interaction with the Galactic tidal field. Indeed, two well defined and extended tidal tails associated with Pal14, likely due to an active process of tidal stripping, have been recently detected (S11). However, no measurement of the degree of mass segregation was available to date and the estimated half-mass relaxation time ($t_{rh}\\sim 20$ Gyr; S11) is much longer than the cluster age \\citep[$t\\sim 10.5$ Gyr;][]{dot08}. On the basis of these results and using N-body simulations, \\citet{zonoo11} concluded that either a primordial mass segregation or a non-canonical initial mass function must have been established in this cluster after the initial gas expulsion.  With the aim of clarifying the dynamical state of Pal14, here we study the radial distribution of its BSS population. In Sect. 2 the adopted photometric dataset is presented. In Section 3 and 4 the BSS and the reference populations are defined and their radial distribution is derived and analyzed. In Section 5 we present our conclusions. ",
        "conclusions": "The radial distribution of BSS in Pal14 is indistinguishable from that of its normal (and less massive) stars. As in the case of the only two other clusters showing the same feature \\citep[$\\omega$Cen and NGC2419;][respectively]{fe06,da08}, this is an observational proof that Pal14 is dynamically young, still far from having established energy equipartition even in its innermost regions. This is in agreement with the extremely long half-mass radius relaxation time ($\\sim 20$ Gyr) recently estimated by S11.  Moreover, our results suggest that the unusually flat mass function measured by \\citet{jo09} cannot be explained by energy equipartition developed during the cluster dynamical evolution, but should be primordial \\citep[as suggested by][]{zonoo11}. We note that since the mass range covered by that study is quite limited, further investigation is needed.  Finally, the negligible degree of relaxation of Pal14 suggests that the observed tidal tails should not  be preferentially populated by low mass stars evaporated from the cluster.  The flat BSS radial distribution also suggests that (as expected in such a low density environment) stellar collisions played a minor role in generating these exotica and affecting the binary population. Hence, as in the case of $\\omega$Centauri and NGC 2419, the BSS we are observing likely derive from the evolution of primordial binaries and can be used to get a rough estimate of the fraction of such a population. As a first consideration, we note that the number of BSS normalized to the sampled luminosity \\citep[in units of $10^4 L_\\odot$, see][]{fe95} is $S4_{\\rm BSS}\\simeq 2$ in $\\omega$Centauri and $S4_{\\rm BSS}\\simeq 3$ in NGC 2419. The same ratio in Pal14 rises to 29 (i.e. a value 10 times larger than what found in the two clusters with similar BSS radial distribution). However this value is not that surprising when compared to the field. In fact, as discussed by \\citet{fe06}, the observed BSS specific frequency $N_{\\rm BSS}/N_{\\rm HB}=0.1$ in $\\omega$Centauri turned out to be $\\sim 40$ times lower that what observed in the field ($N_{\\rm BSS}/N_{\\rm HB}=4$) by Preston \\& Sneden (2000).  The value found in Pal14 is $N_{\\rm BSS}/N_{\\rm HB}\\sim 1$, in much better agreement with the above field sample.  Under the hypothesis that all the BSS are originated by primordial binaries, this possibly suggests that the binary fraction in Pal14 (and in the field) might be much higher than in the other two GCs \\citep[in $\\omega$Cen it amounts to $f_{\\rm bin}\\sim 13\\%$;][]{so07a}. A rough estimate of the binary fraction in Palomar 14 can be derived from the correlation between the measured binary fraction and the cluster integrated magnitude found by~\\citet{so10} in a sample of 18 open and low-density globular clusters.  From their Figure 7, and adopting $M_V=-4.95$ (S11) we predict $f_{\\rm bin}\\sim 30-40\\%$ for Palomar 14.  In addition, the comparison between the number of BSS per unit luminosity \\citep[from][]{fe95} and the fraction of binaries \\citep{so07b} measured in a sample of low-density GCs (in which the collisional channel of BSS formation is expected to be negligible) shows that while BSS-poor GCs (with $8<S4_{\\rm BSS}<13$) host a small fraction of binaries ($9<f_{\\rm bin}<16\\%$), the only cluster (Palomar 12) with a value of $S4_{\\rm BSS}$ similar to Pal14 has a binary fraction of $\\sim 40\\%$.  Despite the uncertainty affecting these estimates, such a result and, even more robustly, the existence of a non-collisional BSS population, suggest that Pal14 might have a non-negligible fraction (of the order of $\\sim 30-40\\%$) of binaries.  Deep and high-quality imaging of Pal14 main sequence are urged for a direct estimate of the binary fraction and a measurement of the mass function in a wide range of masses. Additional spectroscopic campaigns able to more precisely estimate the cluster velocity dispersion would also be very useful to constrain more tightly the dynamical state of the cluster."
    },
    "1107/1107.5606.txt": {
        "abstract": "{ Based on high-resolution ($R\\approx42\\,000$ to $48\\,000$) and high signal-to-noise ($S/N\\approx50$ to 150) spectra obtained with UVES/VLT, we present detailed elemental abundances (O, Na, Mg, Al, Si, Ca, Ti, Cr, Fe, Ni, Zn, Y, and Ba) and stellar ages for 12 new microlensed dwarf and subgiant stars in the Galactic bulge. Including previous microlensing events, the sample of homogeneously analysed bulge dwarfs has now grown to 26.  The analysis is based on equivalent width measurements and standard 1-D LTE MARCS model stellar atmospheres. We also present NLTE Li abundances based on line synthesis of the $^7$Li line at 670.8\\,nm. The results from the 26 microlensed dwarf and subgiant stars show that the bulge metallicity distribution (MDF) is double-peaked; one peak at  $\\rm [Fe/H]\\approx -0.6$ and one at $\\rm [Fe/H]\\approx +0.3$, and with a dearth of stars around solar metallicity. This is in contrast to the MDF derived from red giants in Baade's window, which peaks at this exact value. A simple significance test shows that it is extremely unlikely to have such a gap in the microlensed dwarf star MDF if the dwarf stars are drawn from the giant star MDF. To resolve this issue we discuss several possibilities, but we can not settle on a conclusive solution for the observed differences. We further find that the metal-poor bulge dwarf stars are predominantly old with ages greater than 10\\,Gyr, while the metal-rich bulge dwarf stars show a wide range of ages. The metal-poor bulge sample is very similar to the Galactic thick disk in terms of average metallicity, elemental abundance trends, and stellar ages. Speculatively, the metal-rich bulge population might be the manifestation of the inner thin disk. If so, the two bulge populations could support the recent findings, based on kinematics, that there are no signatures of a classical bulge and that the Milky Way is a pure-disk galaxy. Also, recent claims of a flat IMF in the bulge based on the MDF of giant stars may have to be revised based on the MDF and abundance trends probed by our microlensed dwarf stars. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.5165_arXiv.txt": {
        "abstract": " ",
        "introduction": "Over the coming decade we will see the completion of a large number of major astronomical surveys covering all wavelengths across large swathes of the sky. Such surveys are already underway or being planned, and many more will undoubtedly be initiated in the coming years. Here I focus on the near-infrared surveys currently being carried out by the ESO-Visible and Infrared Survey Telescope for Astronomy (VISTA) at Paranal, Chile, in particular the second deepest tier of the extragalactic public surveys to be undertaken on this telescope, namely the VISTA Deep Extragalactic Observations (VIDEO) survey. I also briefly discuss the various other multi-wavelength surveys that will be carried out over these regions of sky.  I also discuss two of the surveys underway with the {\\sl Herschel Space Observatory}, {\\sl Herschel}-ATLAS and HerMES, which when combined with shorter wavelength data will allow us to trace the obscured activity in the Universe.  Finally, I discuss how the South African Large Telescope and the MeerKAT can play a major role in advancing our understanding of the extragalactic Universe when combined with the imaging surveys underway elsewhere. ",
        "conclusions": ""
    },
    "1107/1107.5812_arXiv.txt": {
        "abstract": "Rapid variability on a time scale much faster than the light-crossing time of the central supermassive black hole has been seen in TeV emission from the blazar PKS 2155-304. The most plausible explanation of this puzzling observation is that the radiating fluid in the relativistic jet is divided into a large number of sub-regions which move in random directions with relative Lorentz factors $\\approx \\Gj$. The random motions introduce new relativistic effects, over and above those due to the overall mean bulk Lorentz factor $\\Gb$ of the jet.  We consider two versions of this ``jets in a jet'' model.  In the first, the ``subjets'' model, stationary regions in the mean jet frame emit relativistic subjets that produce the observed radiation. The variability time scale is determined by the size of the sub-regions in the mean jet frame.  This model, which is motivated by magnetic reconnection, has great difficulty explaining the observations in PKS 2155-304.  In the alternate ``turbulence'' model, various sub-regions move relativistically in random directions and the variability time scale is determined by the size of these regions in their own comoving frames.  This model fits the data much more comfortably.  Details such as what generates the turbulent motion, how particles are heated, and what the radiation process is, remain to be worked out.  We consider collisions between TeV photons emitted from different sub-regions and find that, in both the subjets and turbulence models, the mean bulk Lorentz factor $\\Gb$ of the jet needs to be $>25$ to avoid the pair catastrophe. ",
        "introduction": "On July 28, 2006, the high-frequency peaked BL Lac source PKS 2155-304 had a strong flare in the TeV band, with an average flux that was more than 10 times larger than the typical flux seen at other times \\citep{Aharonian+07}. During this flare, which lasted for about an hour, rapid variability on a time scale $\\approx 300$\\,s was observed. With a galaxy bulge luminosity $M_R = -24.4 $ \\citep{Kotilainen+98}, the expected mass of the black hole (BH) in the nucleus of the galaxy is $M_{\\rm BH} \\approx 1-2 \\times 10^9 M_\\odot$ \\citep{Bettoni+03}. This corresponds to a gravitational radius $R_g \\equiv 2GM/c^2 \\approx 3-6 \\times 10^{14}$\\,cm, and a corresponding time scale $T_g=R_g/c \\approx1-2 \\times 10^4$\\,s.  Amazingly, $T_g$ is nearly two orders of magnitude longer than the variability time.  It is generally accepted that TeV emission in blazars comes from relativistic jets that are pointed towards the observer \\citep{Hinton_Hofmann09}. Opacity arguments suggest that the bulk Lorentz factor of the jet in PKS 2155-304 must be very large, $\\Gb>50$ \\citep{Begelman+08}. At first sight, it might appear that this large Lorentz factor will also explain the rapid variability observed in the source.  However, while relativity can certainly cause the observed variability time to be shorter than $R/c$, where $R$ is the size of the emitting region, there is no simple way for the variability time to be shorter than the gravitational time scale $T_g$ of the central engine. In PKS 2155-304, the variability is faster by a factor of several tens.  The only way to understand the rapid variability in PKS 2155-304 is if one of the following conditions holds: (i) the entire source (engine and emitting region) moves towards the observer with a Lorentz factor $\\approx 50$, or (ii) the emitting region alone moves rapidly towards the observer, and the variability is caused by some local instability in the radiating gas which is insensitive to any time scale associated with the central BH engine, or (iii) the supermassive BH is $\\approx 50$ times less massive than we estimate.  The first option is obviously impossible. We do not expect a $10^9 M_\\odot$ BH to move with a Lorentz factor of 50.  The third option is also unlikely since the BH mass estimate is obtained via the well-known (BH mass)$-$(bulge luminosity) relation \\citep{Magorrian+98}, which has an uncertainty of no more than a factor of a few\\footnote{An exception would arise, of course, if the system consists of a binary black hole and the smaller BH produces the flare \\citep{VolpeRieger}.}.  We are thus led to the second option, but even this is highly non-trivial.  The size of the emitting region is surely at least as large as $R_g$, which means, given the short variability time scale, that only a small fraction of the emitting volume can be involved in the variability. Yet, the large amplitude of the observed fluctuations suggests that the whole emitting region must contribute to the flare. There is an inherent paradox in these conflicting indications.  Building on ideas developed earlier to explain variability in gamma-ray bursts (GRBs, see \\citealt{Lyutikov_Blandford03,Lyutikov06,Narayan_Kumar09,Lazar+09}), Giannios et al. (2009; see also \\citealt{Giannios+10,Nalewajko+11}) proposed a ``jets in a jet'' model as a possible explanation of the TeV flare observed in PKS 2155-304.  According to this model, a relativistic jet moves with an overall bulk Lorentz factor $\\Gb$ towards the observer, but the TeV emission is produced in emitting regions that themseves move relativistically with respect to the mean frame of the jet.  As shown in the context of GRBs \\citep{Narayan_Kumar09,Lazar+09}, such a model naturally produces large amplitude variability on time scales much shorter than the light-crossing time of the emitting volume.  In this work we explore the ``jets in a jet'' model and obtain several constraints that must be satisfied for this model to operate.  Our analysis follows the work of \\cite{Narayan_Kumar09} and \\cite{Lazar+09} on the equivalent problem for GRBs.  In \\S 2 we summarize the relevant observations of PKS 2155-304. In \\S 3 we discuss the overall issue of time scales in relativistic jets. We then describe the ``jets in a jet\" model and its two basic variants --- ``subjets\" (\\S3.1), ``turbulence\" (\\S3.2) --- and analyze the conditions under which each can fit the observations. Both variants of the model involve a large number of emitting regions moving in different directions. Photons emitted in one region can interact with those emitted in another, and if the optical depth is sufficiently large, colliding high energy photons will pair produce and make the source opaque. We examine the optical depth to this process in \\S 5 and consider how the constraint $\\Gb>50$ obtained by \\citet{Begelman+08} is modified in the ``jets in a jet\" scenario.  In \\S 6 we summarize our results and consider the implications. Additional details are given in two Appendices. ",
        "conclusions": "\\label{implications}   The very rapid variability observed in the TeV flux of PKS 2155-304 \\citep{Aharonian+07}, which is faster by a factor of 50 than the gravitational time scale of the central BH, is a remarkable puzzle. Normally, even assuming a relativistic outflow, one does not expect to see variability faster than the gravitational time of the driving engine.  PKS 2155-304 violates this expectation by a large factor.  A possible way out --- apparently the only way out --- is to invoke some version of the ``jets in a jet'' model.  This model was discussed earlier (though this name was not used) in connection with the variability of GRBs \\citep{Lyutikov_Blandford03,Lyutikov06,Narayan_Kumar09,Lazar+09}). The same idea has been shown to explain the variability in PKS 2155-304 \\citep{Giannios+09}.  In the ``jets in a jet\" model the fluctuating pulses in the observed radiation are not produced by the entire jet but by small sub-regions within the jet. These sub-regions move relativistically with respect to the mean frame of the jet. As a result each sub-region produces a beam of radiation that is focused tightly in the direction of its motion.  Only a few sub-regions radiate towards the observer, but these appear anomalously bright as a result of relativistic beaming. Thus, the model explains both the large apparent luminosity and the rapid variability, without violating any energy or light-crossing-time constraints.  We have considered in this paper two versions of the ``jets in a jet'' model. The first version, the ``subjets'' model, assumes that something like magnetic reconnection takes place in many different sub-regions within the jet fluid and that each reconnection site ejects twin relativistic subjets in opposite directions along the local recombining magnetic field. This model is characteristic of a class of models in which the time scale of pulses in the observed light curve is determined by a process that is at rest in the frame of the jet. The role of the relativistic subjet is merely to provide a luminosity boost through beaming.  The second version, the ``turbulence'' model, assumes that some instability in the jet fluid leads to highly relativistic random motions within the medium and as a result the radiation from each sub-region is narrowly beamed in the direction of the local motion. This model is characteristic of a class of models in which the observed variability time scale is determined by a process that takes place in the frame of the relativistically moving sub-region.  Thus, both the observed variability time scale and the luminosity are affected by the motion of the sub-region (in contrast to the subjets model where only the luminosity is modified).  These two models are quite similar in spirit, but they differ in details which become important in the case of extreme objects such as PKS 2155-304.  In the case of the subjets model, using our canonical observational constraints from PKS 2155-304, viz., flare duration $T \\approx 20000$\\,s, variability time $\\delta t \\approx 300$\\,s, pulse duty cycle $\\xi \\approx 0.4$, reconnection inflow speed $\\beta' \\approx 0.1$, we find that the model requires the Lorentz factor of the subjets, as measured in the mean jet frame, to be $\\Gj \\approx 2500$. Given that $\\Gj \\sim \\sqrt{\\sigma}$, where $\\sigma$ is the magnetization parameter of the jet fluid, this value seems unreasonably large.  Perhaps of even greater concern, the model requires the total number of independent reconnection sites, or sub-regions, to be $n_{\\rm tot} \\approx 10^9$.  Both requirements become less stringent if we assume extreme values for the time scales, e.g., $T \\approx 4000$\\,s (i.e., setting $T$ equal to the duration of the observations, assuming that the flare started exactly when the observations began) and $\\delta t \\approx 600$\\,s (the maximum duration of each pulse, see \\citealt{Aharonian+07}). For this choice, $T/\\delta t \\approx 7$, $\\xi \\approx 0.8$, and we find $\\Gj \\approx 170$, $n_{\\rm tot} \\approx 10^6$.  The crux of the problem in the subjects model is the relatively low fiducial value we assume for $\\beta'$ (based on the work of \\citealt{Lyubarsky05}).  The velocity $\\beta' c$ is the speed with which magnetized fluid moves in towards the reconnection point. Since this velocity determines the variability time, a low value of $\\beta'$ implies a correspondingly small size $\\lp$ of the reconnection cell. To compensate, $\\Gj$ has to be very large so that the small amount of energy available within one reconnection cell is strongly beamed to give the luminosity observed in a single pulse of the light curve.  It is possible that the arguments leading to $\\beta' \\approx 0.1$ could be circumvented, allowing a reconnection inflow speed close to $c$. It is also possible that some (unknown) mechanism other than reconnection produces the subjets and that this mechanism transports energy to the subjets at the speed of light. However, even if we set $\\beta'=1$, this is still only marginally acceptable. Using fiducial values for the other parameters we find $\\Gj \\approx 80$, $n_{\\rm tot} \\approx 10^6$. We obtain a reasonable solution only if we choose both a high value of $\\beta' \\approx 1$ and extreme values of the time scales, $T \\approx 4000$\\,s, $\\delta t \\approx 600$\\,s.  In this case, we find $\\Gj \\approx 5$, $n_{\\rm tot} \\approx 10^3$.  Or, we could choose $\\beta' \\approx 1$ and instead of changing $T$ and $\\delta t$, we could select a smaller BH mass, $M_9 \\approx 0.5$, which gives $\\Gj\\approx 20$, $n_{\\rm tot}\\approx 4\\times 10^4$.  We thus conclude that the subjets model works only if magnetic reconnection in relativistic plasmas proceeds at the speed of light,  much faster than currently thought.  This alone may not be sufficient to explain PKS 2155-304.  We may also need to push the observational estimates of time scales or BH mass to extreme values.  The situation in the case of the turbulence model is quite different. In contrast to the subjets model, here, even with fiducial values of parameters, the implied conditions in the jet are quite reasonable. For instance, with $T \\approx 20000$\\,s, $\\delta t \\approx 300$\\,s, $\\xi \\approx 0.4$, we find $\\Gj\\approx5-7$, $n_{\\rm tot}\\approx10^3$.  Such conditions appear quite reasonable for a relativistically moving high-energy source.  The weakness of the turbulence model is that it is short on details. We do not have a physical model of what causes the turbulence and what the limits of this process are. (This is in contrast to the subjets model where we have a very specific process in mind -- reconnection -- which immediately gives a physical constraint $\\beta'\\sim0.1$.)  In fact, the word ``turbulence'' itself is merely a code for random motions. The actual dynamics may not be truly turbulent in the sense it is normally understood.  Certainly, we cannot have hydrodynamic turbulence since motions with $\\Gj\\approx 5$ are highly supersonic with respect to the maximum allowed sound speed, $c_{\\rm s,max}=c/\\sqrt{3}$.  In principle, MHD turbulence is compatible with the model since wave speeds can reach up to Lorentz factors $\\Gamma_{\\rm wave}\\sim\\sqrt{\\sigma}$, where $\\sigma$ is the magnetization parameter.  Therefore, with a sufficiently strongly magnetized medium ($\\sigma \\gg 1$), random motions with $\\Gj \\approx 5-7$ could be supported.  The heating of particles and the production of radiation is also unexplained in the turbulence model.  In the subjets model, reconnection automatically produces relativistic beams of particles and it is not hard to imagine that these particles will radiate in the ambient magnetic field.  (Note, however, that the dissipation efficiency in reconnection may be rather low unless the reconnecting fields coming in from the two sides are practically parallel, see \\citealt{Lyubarsky05}).  In the case of the turbulence model, the most obvious candidate for particle heating is a shock, but this is unlikely to work since highly magnetized shocks are very inefficient at particle heating \\citep{Kennel_Coroniti_84,Narayan+11}.  Thus, it may be necessary to invoke heating through reconnection inside the turbulent blobs (which would immediately introduce inefficiency through a $\\beta'$-like factor, just as in the subjets model). Perhaps the most promising idea is that the particles are heated directly by waves in the turbulent magnetized medium, but the details remain to be worked out.  Summarizing, we find the situation far from clear. The subjets model with magnetic reconnection provides a natural way to produce the required ``jets within a jet.'' However, given our current understanding of relativistic reconnection \\citep{Lyubarsky05}, viz., that the inflow speed towards the reconnection site can be no larger than a tenth the speed of light, this model has great difficulty explaining the observations in PKS 2155-304.  The turbulence model \\citep{Narayan_Kumar09,Lazar+09} apparently has no difficulty fitting the data.  However, all we have is a broad outline of this model, and there is no physical picture of how the ``turbulence'' in this model is produced or how this turbulence heats particles and produces the observed radiation.  In \\S4 we analyze the pair production opacity for TeV photons in the ``jets in a jet'' model.  We show that collisions of beams emitted from different sub-regions are common and hence the opacity due to these beam-beam collisions is more important than the opacity within a single emitting sub-region.  Allowing for this effect we find that, regardless of which version of the ``jets in a jet\" model we consider, the minimum bulk Lorentz factor $\\Gb$ of the jet is $\\approx 25$, which is lower than the limit $\\approx 50$ obtained by \\citet{Begelman+08}.  Although the reduction in the value of $\\Gb$ is only a factor of 2, it is nevertheless a significant revision since the new value is closer to the typical bulk Lorentz factors $\\approx10-20$ found in relativistic jets in other blazars.  This work was supported in part by NSF grant AST-1041590 and NASA grant NNX11AE16G (RN), and by the Israel Center for Excellence for High Energy Astrophysics and an ERC advanced research grant (TP).  \\appendix \\numberwithin{equation}{section}"
    },
    "1107/1107.2421_arXiv.txt": {
        "abstract": "At a density near a few $\\times 10^7$ g cm$^{-3}$, the subsonic burning in a Type Ia supernova enters the distributed regime (high Karlovitz number).  In this regime, turbulence disrupts the internal structure of the flame, and so the idea of laminar burning propagated by conduction is no longer valid.  The nature of the burning in this distributed regime depends on the turbulent Damk\\\"ohler number ($\\Da_T$), which steadily declines from much greater than one to less that one as the density decreases to a few $\\times 10^6$ g cm$^{-3}$. Classical scaling arguments predict that the turbulent flame speed $s_T$, normalized by the turbulent intensity $\\check{u}$, follows $s_T/\\check{u}=\\Da_T^{1/2}$ for $\\Da_T\\ltaprx 1 $.  The flame in this regime is a single turbulently-broadened structure that moves at a steady speed, and has a width larger than the integral scale of the turbulence.  The scaling is predicted to break down at $\\Da_T\\approx1$, and the flame burns as a turbulently-broadened effective unity Lewis number flame.  This flame burns locally with speed $s_\\lambda$ and width $l_\\lambda$, and we refer to this kind of flame as a $\\lambda$-flame.  The burning becomes a collection of $\\lambda$-flames spread over a region approximately the size of the integral scale.  While the total burning rate continues to have a well-defined average, $s_T \\sim \\check{u}$, the burning is unsteady.  We present a theoretical framework, supported by both 1D and 3D numerical simulations, for the burning in these two regimes.  Our results indicate that the average value of $s_T$ can actually be roughly twice $\\check{u}$ for $\\Da_T\\gtaprx 1$, and that localized excursions to as much as five times $\\check{u}$ can occur. We also explore the properties of the individual flames, which could be sites for a transition to detonation when $\\Da_T\\sim1$. The $\\lambda$-flame speed and width can be predicted based on the turbulence in the star (specifically the energy dissipation rate $\\varepsilon^*$) and the turbulent nuclear burning time scale of the fuel $\\tau_{\\mathrm{nuc}}^T$.  We propose a practical method for measuring $s_\\lambda$ and $l_\\lambda$ based on the scaling relations and small-scale computationally-inexpensive simulations.  This suggests that a simple turbulent flame model can be easily constructed suitable for large-scale distributed supernovae flames.  These results will be useful both for characterizing the deflagration speed in larger full-star simulations, where the flame cannot be resolved, and for predicting when detonation occurs. ",
        "introduction": "\\label{sec:intro}  In \\citet{Aspden08a} (henceforth Paper~I) three-dimensional simulations of resolved flames were performed that examined the interactions between turbulence and a carbon-burning flame in a Type Ia supernova at different densities.  Because of the strong dependence of flame width and speed on density, this study can be viewed as a survey of the behavior of the flame at variable Karlovitz number, \\begin{equation} \\Ka=\\sqrt{\\frac{\\check{u}^3}{s_L^3}\\frac{l_L}{l}}. \\end{equation} Here $s_L$ and $l_L$ are the laminar flame speed and width, respectively, and $\\check{u}$ and $l$ are the turbulent intensity (rms velocity) and integral length scale (defined conventionally as the integral of the longitudinal correlation function). For $\\Ka\\ltaprx1$, the flame is laminar with propagation determined by a balance between conductive and burning time scales. Such laminar flames have a large Lewis number, which is to say thermal diffusion occurs much faster than carbon diffusion.  This means that perturbing the flame surface into the ash leads to a focusing of heat by diffusion, enhancing the burning rate, burning away the perturbation. Similarly, perturbing the flame surface into the fuel leads to a defocusing of heat by diffusion, decreasing the burning rate.  As a result, the flames are thermodiffusively stable. At small-to-moderate Karlovitz numbers ($\\Ka\\ltaprx 1$), it was found that this thermodiffusively-stable nature led to a balance between local enhancement of the flame and local extinction, and so the turbulent flame speed remained close to the laminar value.  However, once the Karlovitz number was sufficiently large ($\\Ka \\gtaprx 10$), turbulence was sufficiently strong that the turbulent mixing dominated thermal diffusion, and the flame became completely stirred (it resembled a turbulent mixing zone) and its width was greatly broadened. The turbulent flame width was also much greater than the integral length scale.  While the local burning rate was greatly reduced, the overall flame speed was a factor of five or six times the laminar flame speed due to the enhanced volume of the burning.  A key result from Paper~I was that at high Karlovitz number ($\\Ka\\approx230$), turbulence dominates the mixing of fuel and heat while thermal diffusion plays only a minor role. Figure \\ref{Fig:PdfFuelTemp} shows a joint probability density function (\\jpdf) of fuel and temperature from Paper~I. The solid red line shows the distribution from the flat laminar flame where thermal diffusion is the dominant mixing process, and the solid black line shows the distribution of fuel burning isobarically with no thermal or species diffusion.  The agreement with the \\jpdf\\ demonstrates that the turbulent flame is burning at an effective Lewis number close to unity - the mixing is chiefly due to turbulence.  The simulations in Paper~I were performed in small domains to ensure that the laminar flame was fully-resolved -- the domain size was approximately twenty-five times the laminar flame width.  The aim of this paper is to investigate turbulent flame speeds in larger domains, with the aim of predicting high Karlovitz flame speeds in a full star, and specifically testing predictions made in Paper~I.  In this paper, we will refine these predictions by giving an in-depth theoretical description of the distributed burning regime, which are then compared with one- and three-dimensional simulations. ",
        "conclusions": "\\label{sec:conclusions}  New one- and three-dimensional simulations have been presented to clarify and quantify the nature of carbon burning in a Type Ia supernova in the distributed burning regime. The characteristics of distributed burning depend upon the Damk\\\"ohler number, and a range of $\\Da_T$ from 0.006 to 26 has been explored.  For $\\Da \\ltaprx 1$, the expected scaling relations were demonstrated, \\begin{equation} \\frac{s_T}{\\check{u}} = \\Da_T^{1/2}, \\quad\\mathrm{and}\\quad \\frac{l_T}{l} = \\Da_T^{-1/2}. \\label{eq:concRelations} \\end{equation}  For $\\Da_T \\gtaprx 1$, these relations break down and the turbulent flame reaches a maximum local flame speed $s_\\lambda$ and width $l_\\lambda$, measured on the scale of $l_\\lambda$.  This $\\lambda$-flame interacts with turbulent eddies with length scales between $l_\\lambda$ and the integral length scale  $l$, which enhances the flame surface area.  The average turbulent flame properties are constrained by the large scales  \\citep[see also][]{Woosley09}.  This leads to an average overall turbulent flame speed, which to factor-of-two accuracy is given by the turbulent intensity, $\\check u$. For the range of $\\Da \\gtaprx 1$ studied here, a good approximation is $s_T = 2 \\check u$.  \\subsection{Consequences for Large-Scale Simulations}  Large-scale simulations (i.e.\\ of a full-star) will require a turbulent flame model.  One of the consequences of the results presented here is that the $\\lambda$-flame is ideally suited to the level-set approach, as it burns locally with speed $s_\\lambda$.  This means that a turbulent flame speed can be evaluated in the following manner.  Assume that the turbulence is sufficiently large and intense, i.e.\\ in the stirred-flame regime ($\\Da_T>1$ and $\\Ka\\gtaprx10$). Then the turbulence can be characterized by the energy dissipation rate $\\varepsilon^*=\\check{u}^3/l$, and the burning depends on the turbulent nuclear burning time scale $\\tau_{\\mathrm{nuc}}^T$.  The nuclear burning time scale is equal to $l_\\lambda/s_\\lambda$ and the turbulent intensity at the length scale of the $\\lambda$-flame is $s_\\lambda$ and so $\\varepsilon^*=s_\\lambda^3/l_\\lambda$.  Solving for $s_\\lambda$ and $l_\\lambda$ gives \\begin{equation} s_\\lambda = \\sqrt{\\varepsilon^*\\tau_{\\mathrm{nuc}}^T}, \\quad\\mathrm{and}\\quad l_\\lambda = \\sqrt{\\varepsilon^*{\\tau_{\\mathrm{nuc}}^T}^3}. \\label{eq:lambda} \\end{equation} Note these expressions do not require a fixed Karlovitz number (it's incorporated by $\\varepsilon^*$), and the the length scale is the same as predicted in Paper I.  To implement a level-set approach for a $\\lambda$-flame, $\\tau_{\\mathrm{nuc}}^T$ should be evaluated as a function of fuel density and temperature (using the approach outlined in section \\ref{sec:wsr} or \\cite{Woosley09}) and coupled with the (local) energy dissipation rate to evaluate the $\\lambda$-flame speed according to equation (\\ref{eq:lambda}).  If the resolution requirements of Paper~I were applied to the $\\lambda$-flame recovered in the present paper (i.e.\\ $l_\\lambda\\approx4\\Delta x$ with $256^2$ cells in cross-section), an integral length scale could be achieved that is of order thousands of times larger than the original study.  \\subsection{Consequences for Detonation}  The instantaneous flame speed for $\\Da_T \\gtaprx 1$ is irregular, with frequent excursions up to 2.5 times the average, which is approximately twice the turbulent intensity for the simulations presented here. These variations could be crucial if a transition to detonation is to occur. Detonation will not happen in the laminar regime, i.e.\\ at high density, because each flamelet burns at a steady rate and has a thickness that is far below the critical mass for detonation. It is also unlikely to happen in the well-stirred reactor regime, because the burning is steady and slower than the turbulence on the integral scale, which in turn is subsonic.  The burning time scale is very long in the well-stirred regime, making it very difficult to burn, for example, a 10 km region on a sound crossing time.  Thus if detonation is to occur spontaneously, it is likely to happen in the stirred-flame regime.  Even there though, $\\Da_T$ should not be too large, or the mixed regions will  not burn supersonically (recall $s_T\\rightarrow\\check{u}$ for large $\\Da_T$). The best conditions therefore occur where there are multiple $\\lambda$-flames across an integral length scale, i.e.\\ for small $\\Da_T \\gtaprx 1$ \\citep{Woosley07,Woosley09}. The calibration of LEM determined by comparison with the three-dimensional studies suggest a normalization constant of $C = 5$ is most appropriate for $\\Da > 1$ and this is the value used by \\citet{Woosley09}. This strengthens their conclusion that a detonation is possible if $\\check{u}$ exceeds about one-fifth sonic."
    },
    "1107/1107.5940_arXiv.txt": {
        "abstract": "We analyse the chromospherical activity of stars with extrasolar planets and search for a possible correlation between the equivalent width of the core of Ca\\,II\\,K line and orbital parameters of the planet. We found a statistically significant evidence that the equivalent width of the Ca\\,II\\,K line reversal, which originates in the stellar chromosphere depends on the orbital period $P_{\\mathrm{orb}}$ of the exoplanet. Planets orbiting stars with $T_{\\mathrm{eff}}<5\\,500$\\,K and with $P_{\\mathrm{orb}}<20$\\,days generally have much stronger emission than planets at similar temperatures but at longer orbital periods. $P_{\\mathrm{orb}}=20$\\,days marks a sudden change in behaviour, which might be associated with a qualitative change in the star-planet interaction. ",
        "introduction": "The question of possible existence of star-planet interaction is currently studied in many ways. Based on the observations in the optical region \\cite[Shkolnik et al. (2005, 2008)]{Shkolnik08} discovered the planetary induced variability in the cores of Ca II H \\& K, H$\\alpha$ and Ca\\,II IR triplet in a few planet hosting stars. \\cite[Knutson et al. (2010)]{Knutson10} found a correlation between the chromospheric activity of the star and presence of the stratosphere on the planet. Consequently, \\cite[Hartman (2010)]{Hartman10} found a correlation between the surface gravity of Hot Jupiters and the stellar activity. Recently \\cite[Canto Martins et al. (2011)]{Martins11} searched for correlation between planetary parameters and the $\\log R'_{\\mathrm{HK}}$ parameter but didn't reveal any convincing proof for such a phenomenon. ",
        "conclusions": ""
    },
    "1107/1107.0562_arXiv.txt": {
        "abstract": "We discuss theoretical AGB predictions for hydrogen-deficient PG 1159 stars and Sakurai's object, which show peculiar enhancements in He, C and O, and how these enhancements may be understood in the framework of a very late thermal pulse nucleosynthetic event. We then discuss the nucleosynthesis origin of rare subclasses of presolar grains extracted from carbonaceous meteorites, the SiC AB grains showing low  $^{12}$C/$^{13}$C in the range  2 to 10 and the very few high-density graphite grains with $^{12}$C/$^{13}$C around 10. ",
        "introduction": "\\subsection{PG 1159 stars: extremely hot Post-AGB stars}  The hydrogen-deficiency in extremely hot post-AGB stars of spectral class PG 1159, which includes about 40 stars, with  $T_{\\rm eff}$ ranging from 75\\,000 and 200\\,000 K and log $g$ from 5.5 to 7.5, is probably caused by a very late thermal pulse in the He shell (VLTP, Sch{\\\"o}nberner 1979, Iben 1984) while the post-AGB star is in the hot WD cooling sequence. Because of the high $T$ $_{\\rm eff}$ in PG 1159 stars, all species are highly ionized and, hence, most metals are only accessible by UV spectroscopy. A passionate work has been conducted in the last 20 years by Klaus Werner and collaborators (Werner \\& Herwig 2006, Werner et al. 2009, 2010 and references therein). In Table 1 we report in particular the range of peculiar abundances of He, C, N, O estimated in PG 1159 stars, where the mass fraction of He ranges between 0.30  and 0.85, C between 0.15 and 0.40, N between 0.001 and 0.01, and O between 0.02 and 0.2. The energy released by the VLTP forces the stellar radius to inflate and the star to cool and  proceed back toward a born-again AGB. At the maximum extension of the convective thermal instability the very small residual and inactivated H shell is likely engulfed by the pulse and severly depleted, so that the usually hidden He-, C-rich and s-rich He intershell is eventually exposed to the photosphere. PG 1159 stars are seemingly descendants of [WC] stars, which show similar HeCNO peculiarities.  \\subsection{Sakurai's object (V4334 Sgr)} \\begin{figure} \\includegraphics[angle=-90,width=6.5cm]{gallino_fig1.eps} \\includegraphics[angle=-90,width=6.5cm]{gallino_fig2.eps} \\caption{Left panel:  Sakurai's object spectroscopic data by Asplund et al. (1998) compared with the [El/Fe] distribution in the He intershell for different AGB models of $M^{\\rm AGB}_{ini}$ = 3 $M_\\odot$, [Fe/H] = $-$0.3, case ST/4.5 at the last thermal pulse, or for [Fe/H] = $-$0.5 and case ST/12. Right panel: FG Sge spectroscopic data by Gonzalez et al. (1998) compared with  an AGB model of $M^{\\rm AGB}_{ini}$ = 1.5 $M_\\odot$, [Fe/H] = $-$0.3, case ST $\\times$ 1.3 at the last thermal pulse. The lower curve is for the distribution in the envelope at the last thermal pulse for the same AGB model (see text). } \\label{fig1} \\end{figure}  Sakurai's object (V4334 Sgr) was discovered in 1996 and soon recognized to be the central star of an old planetary nebula (6000 yr ago). It recently underwent a VLTP and is now a hydrogen-deficient born-again AGB star. The estimated mass fractions He/C/O = 0.90/0.07/0.03 and other element abundances (Asplund et al. 1998, their Figure 6) are reported in Table 1 and compared with PG1159 stars. This composition is based on the choice of the best adopted atmospheric model with  C/He $\\approx$ 0.10 (by number). Note that there is a \"carbon problem\" for Sakurai's object similar to that encountered for the hydrogen-deficient RCrB stars. Indeed, the spectroscopic C abundance derived from CI lines is about 0.6 dex smaller than the one deduced from the selected model atmosphere. The same is the case for [Fe/H]. However, the relative abundance ratios, like [El/Fe], are scarcely dependent on the C/He choice. In Fig. 1 left panel we compare the observed [El/Fe] data of Sakurai's object with predicted abundances in the He intershell at the last thermal pulse for an AGB model of $M^{\\rm AGB}_{ini}$ = 3 $M_\\odot$, [Fe/H] = $-$0.3, case ST/4.5. Given the uncertainty of the initial metallicity, we also plot in the figure a similar predicted [El/Fe] distribution for $M^{\\rm AGB}_{ini}$ = 3 $M_\\odot$, [Fe/H] = $-$0.5 and case ST/12 (dashed line). In this case, Sc and Rb appear better reproduced, but the reverse is true for C and Cu. As to Sr, its abundance in October 1996 was overestimated. The presence of  carbon dust buffers around Sakurai's object may in general introduce a noticeable uncertainty in spectroscopic abundances. V 605 Aql (Nova Aql 1919) is a second star having likely  suffered a VLTP about 90 yr ago. Clayton et al. (2006) estimated He/C/O = 0.54/0.40/0.05. Both V106 Aql and Sakurai's object showed peculiar rapid declines and fading characteristic of episodic carbon dust emission, as in the case of R CrB stars. However, as discussed below, several R CrB stars likely originated in a completely different way, as binary WD mergers.  \\subsection{H ingestion and partial burning in the He intershell}  In order to produce a consistent amount of N and to achieve the low $^{12}$C/$^{13}$C observed in Sakurai's object, ingestion and burning of hydrogen in a TP  is impossible when the H shell is still active. It  may work when a very thin H envelope is left after the star leaves the AGB (Herwig et al. 1999, Miller Bertolami et al. 2006). A quite low $^{12}$C/$^{13}$C $\\le$ 10 results while a consistent amount of $^{14}$N is built up. During the AGB phase, standard elemental mass fractions in the He intershell after a thermal pulse are He/C/O = 0.75/0.20/0.005. Higher C and O abundances may be obtained by  including an efficient overshoot at the base of the convective thermal pulse in the TP-AGB phase (Herwig et al. 1997). Alternatively, proper account should be given to the peeling effect by mass loss both at the tip of the AGB and in the early phase of the post-AGB track. There, a \"superwind\" of up to several 10$^{-5}$ $M_\\odot$/yr has been measured. Lawlor \\& MacDonald (2006) introduced these effects in their stellar evolution code in an wide spectrum of initial masses and metallicities. Another important factor is the thickness of the He buffer, which  decreases with increasing the CO core mass, that is with the initial stellar mass. The authors showed that chemical peculiarities observed in stars having suffered the VLTP do not strictly require overshoot in the AGB phase. One should also consider that the bottom of the VLTP is degenerate, different from what occurs during the AGB phase, with the possibility to further increase C and O.   \\begin{table*} \\caption{PG 1159 stars, Sakurai's object, AGB  He-intershell predictions, extra needs. The spectroscopic data for Sakurai's objects are from Asplund et al. (1998), their Fig. 6. As to Ne, our AGB predictions refer to $^{22}$Ne, the most abundant isotope in the He intershell. The range of $^{19}$F for AGB He-intershell predictions increases with the number of thermal pulses, i.e. with the initial AGB mass.}   \\label{table1} \\begin{center} \\resizebox{13cm}{!}{\\begin{tabular}{llllll} \\\\ \\hline &PG 1159 & Sakurai's object & He intershell & extra needs       \\\\ $^{12}$C/$^{13}$C           & not measurable & 2 to 5  & NO ${13}$C   &   H-b \\\\ H                &deficient            &deficient &0.& \\\\ He               & from 0.30 to 0.85   &  0.90     &  0.75 & partial He-b \\\\ C                & from 0.15 to 0.60   &0.07     &  0.2 &partial He-b \\\\ N               & from 0.001 to 0.01        & 0.01 & 0. &          H-b \\\\ O               & from 0.02 to 0.20 & 0.03       &  0.005 & partial He-b \\\\ F               & 1 to 250$\\times$solar  & not detectable       & 1 to 250 & \\\\ Ne          & 0.02             & 0.02& 0.02 & \\\\ Na             &&25$\\times$solar & 8$\\times$solar & \\\\ Si              & from solar to 0.5 $\\times$ solar   & 7$\\times$solar  & solar & \\\\ S           & from solar to 0.1$\\times$solar         &2.5$\\times$solar  &  solar & \\\\ P               & from solar to 2 $\\times$ solar   &          &  solar to 4$\\times$solar & \\\\ Ar          & solar      &                           &  solar & \\\\ Fe              & solar to 0.1 $\\times$ solar & 0.6$\\times$solar  &  solar to 0.7$\\times$solar &  \\\\ Ni              & not enhanced      &   8$\\times$solar  &  not enhanced & \\\\ Cu              && 55$\\times$solar &25$\\times$solar& \\\\ Zn              && 55$\\times$solar &20$\\times$solar& \\\\ s-elements  & impossible to detect  & highly enhanced    & highly enhanced& \\\\ \\hline \\end{tabular}} \\end{center} \\end{table*}   \\subsection{FG Sge } The peculiar FG Sge has also been assumed to have recently suffered a VLTP. Over the last 120 years FG Sge evolving from a hot post-AGB star to a present cool and born-again AGB. Figure 1 (panel right)  shows the FG Sge spectroscopic data by Gonzalez et al. (1998) (full triangles) compared with theoretical predictions (upper curve) in the  He-intershell after the last thermal pulse for an AGB initial mass of 1.5 $M_\\odot$, metallicity [Fe/H] = $-$0.3 and the $^{13}$C pocket choice ST $\\times$ 1.3. Note the huge [ls/Fe], [hs/Fe] and [Pb/Fe], of the order of 3 dex each. The  high value [Eu/Fe] = 2 dex observed is s-process Eu, in agreement with the typical s-process expectation [La/Fe]$_s$ $\\approx$ 1 dex. Moreover, Gonzalez et al. estimated $^{12}$C/$^{13}$C $>$ 10, and provided first spectroscopic evidence of H-deficiency. Jeffery \\& Sch{\\\"o}nberner (2006) reanalyzed all extant spectroscopic data and atmospheric parameters, raising doubts on the huge s-process abundances derived by Gonzalez et al. (1998). The lower  $T_{\\rm eff}$ = 5500 K chosen for the  model atmosphere brought Jeffery \\& Sch{\\\"o}nberner to conclude that the s-process element abundances are more that one order of magnitude less than inferred, most likely inherited already during its previous  AGB phase. They concluded that FG Sge suffered a late thermal pulse (LTP), not a VLTP, then evolving back to a born-again AGB. In the lower curve of Figure 1, we compare the envelope AGB model prediction at the last thermal pulse. Note that the two indicators of the s-process distribution,  [hs/ls] and [Pb/hs], would remain unaltered. The spectroscopic heavy elements have been reduced by 1.5 dex (empty triangles), what corresponds to a typical factor of 30 of dilution of He intershell material mixed with the envelope. Note that predicted [C/Fe] would better compare with a LTP solution.  \\subsection{The R CrB  stars} \\begin{figure} \\includegraphics[angle=0,width=6cm]{gallino_fig3.eps} \\includegraphics[angle=0,width=6cm]{gallino_fig4.eps} \\caption{Permil variation with respect to solar of Ca isotopes of graphite presolar grains g-34 and g-40 compared with  He intershell predictions of two different AGB models (adapted from Jadhav et al. 2008). } \\label{fig2} \\end{figure} So far about 50 R CrB have been discovered. Their atmospheres are extremely hydrogen deficient and carbon rich. Another distinctive feature of some R CrB stars is the enormous F detection, in the range 1,000 to 8,000 times solar for  [Fe/H]  in the range  $-$0.5 to  $-$2.0 (Pandey et al. 2008). Such drastic $^{18}$O and $^{19}$F excesses indicate that the  merging of a CO-WD with a He-WD gives rise to partial He burning and production of $^{18}$O via  $\\alpha$-capture on $^{14}$N, accompanied by $^{18}$O(p,$\\gamma$)$^{19}$F. Detailed nucleosynthesis calculations for these peculiar objects are not easy however. ",
        "conclusions": ""
    },
    "1107/1107.2751_arXiv.txt": {
        "abstract": "{We develop a model of production of the very dense clumps of DM in RD epoch due to the accretion of DM on the loops of cosmic strings as the seeds. At some time the loops disappear, for example due to the gravitational radiation, and the remaining dense clumps produce the enhancement of the annihilation signal. We take into account the velocity distribution of the strings, and consider the two extreme regimes of DM annihilation: fast decay and continuous evaporation. The produced annihilation flux of gamma radiation is detectable, and for some parameters of DM particles and the strings can exceed the extragalactic flux of the gamma-radiation observed by Fermi. For the fixed parameters of DM particles (e.g. neutralino with fixed masses and cross-section of annihilation) one can obtain the limits on the basic string parameter, tension $\\mu$, which is stronger than  (more general) limits obtained from WMAP observations, cosmological nucleosynthesis and gravitational lensing. In particular for the neutralino with 100 GeV mass we exclude the interval $5\\times10^{-10}<G\\mu/c^2<5.1\\times10^{-9}$.} ",
        "introduction": "The linear topological defects --- cosmic strings can be formed in the early cosmological phase transitions (see for a review \\cite{VilSheBbook}, \\cite{Vil05}). Along with the infinite strings there is possibility of closed loops formation in the network of curved cosmic strings due to their interconnections. According to numerical simulations, after a long transient stage a true scaling regime sustained, then the typical distances between the strings and the coherence length both scale in proportion with the horizon scale \\cite{alpha}. A string loop formed at the cosmological time $t_i$ has the length $l\\simeq\\alpha ct_i$, where in the scaling regime $\\alpha\\simeq0.1$ according to \\cite{alpha} and \\cite{alpha2}, although other values of $\\alpha$ were obtained in other works (see for example \\cite{alpha3}, \\cite{alpha4}), down to $\\alpha\\sim10^{-3}$.  A fundamental characteristic of the string is the mass per unit length $\\mu\\equiv M_l/l$ or the tension, which is of the order of symmetry breaking energy squared $\\eta^2$. For example, the grand-unification-scale strings have $G\\mu/c^2\\sim10^{-6}$, where $G$ is the Newton's constant. There are several restrictions on $\\mu$. From CMB observations it follows $G\\mu/c^2\\le2\\times10^{-7}$ \\cite{PogWasWym06}. The bound $G\\mu/c^2\\le10^{-7}$ was obtained from the study of nucleosynthesis \\cite{alpha}. Search for the pairs of galaxies images consistent with the gravitational lensing of the cosmic string presents the limit $G\\mu/c^2<3\\times10^{-7}$ at 95\\% confidence level \\cite{Christiansen08}. In \\cite{OluVil06} the formation of stars in the first dark matter DM haloes seeded by the loops was considered. It was found that $G\\mu/c^2<3\\times10^{-8}$ to avoid collision with WMAP data on the reionization redshift. Searches for the gravitational wave bursts from strings by LIGO provide the joint constraints on the strings parameters ($\\mu$ and interconnection probability) \\cite{LIGO}, but $G\\mu/c^2$ is only weakly restricted in comparison with the constraints above. And finally, the strongest bound $G\\mu/c^2\\le4\\times10^{-9}$ was obtained from the pulsar timing \\cite{RvanHaas11}.  In this work we present new constraint on $\\mu$ which was obtained from the DM particles annihilation in the dense clumps seeded by the loops at the cosmological stage of radiation dominance. The direct detections of DM particles are the promising but still elusive experimental problems, therefore the search of indirect signature of the DM is important for clarifying the DM origin. The promising indirect signature --- DM particles annihilation would proceed more efficiently (it can be boosted by several orders) if the Galactic halo is filled by the dense DM substructures or DM clumps. The cosmic string loops produce very dense clumps due to their early formation. Only low velocity loops can produce the clumps. The probability of the low velocity loop formation is very small, but even tiny fraction of the formed loops may produce the dense clump population and significant annihilation signal.  The clumps formation at the radiation dominated (RD) stage was studied in details by \\cite{KolTka94}. In the particular case of loops' density perturbations the maximum density of clump is  restricted due to adiabatic expansion of the already formed clump after the loop gravitational evaporation. We found the modification of this restriction in the case then the loop decays before the clump virialization. In this case the clumps can reach density $\\rho_{\\rm cl}\\gg140\\rho_{\\rm eq}$, where $\\rho_{\\rm eq}$ is the density at equality. The comparison of the resulting annihilation signal with the Fermi-LAT data allows us to obtain the restriction on the string parameter $\\mu$. It must be pointed out that our constraints were obtained by supposing that the DM can annihilate, and we take the $\\sim100$~GeV neutralino as the most promising particle candidate. The constraints will be different for other DM models. In this sense the obtained constraints must be considered as joint constraints on the properties of the string loops and DM particles. ",
        "conclusions": "\\label{discussionsec}  In this work we explore the possibility of dark matter annihilation in the dense cores of clumps that were seeded by cosmic string loops before the matter-radiation equality. We calculate the evolution of the clumps around evaporating loops. Only the low-velocity loops result in the clumps formation. At the same time, even the low-velocity tail of the loop distribution produces DM clumps with the observable signature in the annihilation products. The adiabatic argument conservation doesn't prevent the formation of clumps with densities $\\rho_{\\rm cl}\\gg140\\rho_{\\rm eq}$, if the decay of the loops occurs before the time of the clumps virialization. Therefore the DM clumps produced by the loops can be the very dense objects with high cumulative luminosity in the gamma-rays.  The combined constraints on the loops parameters ($\\mu$ and the distribution over lengths) and parameters of DM particles were obtained. For the 100~GeV neutralino DM the range $5\\times10^{-10}<G\\mu/c^2<5.1\\times10^{-9}$ was excluded because of the huge gamma-ray annihilation signal above the Fermi-LAT data for these parameters. Along with the preferable neutralino dark matter candidate with the mass 100~GeV, we consider the masses $10$~TeV, $1$~TeV, and $10$~GeV and explore the two limiting assumptions about the loop evolution: fast decay and continuous evaporation approximations. The corresponding constraints on the annihilational cross-section are shown at the figure~\\ref{figannl}. We restrict ourself only by the preferable value $\\alpha=0.1$. The results will change for different $\\alpha$.  In the case of the cosmic superstrings, the reconnection probability can be $\\ll 1$, resulting in the much higher number density of loops. This could lead to even stronger constraints in comparison with the presented in this paper.  The Fermi-LAT data are used as the upper limit. In principle the annihilation of DM in the clumps can explain the observed signal for the particular values, for example $G\\mu/c^2\\simeq5\\times10^{-10}$. The necessity of the DM annihilation can arise if the ordinary astrophysical sources give too small signal in comparison with the observations."
    },
    "1107/1107.4341_arXiv.txt": {
        "abstract": "We present the first variability study of the 1720 MHz OH masers located in the Galactic Center. Most of these masers are associated with the interaction between the supernova remnant Sgr\\,A\\,East and the interstellar medium, but a few masers are associated with the Circumnuclear Disk (CND). The monitoring program covered five epochs and a timescale of 20-195 days, during which no masers disappeared and no new masers appeared. All masers have previously been detected in a single epoch observation about one year prior to the start of the monitoring experiment, implying relatively stable conditions for the 1720~MHz OH masers. No extreme variability was detected.  The masers associated with the northeastern interaction region between the supernova remnant and the $+50$~\\kms\\ molecular cloud show the highest level of variability. This can be explained with the $+50$~\\kms\\ molecular cloud being located behind the supernova remnant and with a region of high OH absorbing column density along the line of sight. Possibly the supernova remnant provides additional turbulence to the gas in this region, through which the maser emission must travel.  The masers in the southern interaction region are located on the outermost edge of Sgr\\,A\\,East which line of sight is not covered by either absorbing OH gas or a supernova remnant, in agreement with the much lower variability level observed. Similarly, the masers associated with the CND show little variability, consistent with them arising through collisions between relatively large clumps of gas in the CND and no significant amount of turbulent gas along the line of sight. ",
        "introduction": "Masers observed in the 1720 MHz satellite line of OH are considered signposts of regions of shocked gas, since they often are associated with supernova remnants (SNRs) \\citep[e.g.,][]{frail94,green97}. The 1720\\,MHz masers are not exclusively found in SNRs but are also detected in star forming regions \\citep[e.g.,][]{macleod97,szymczak04,niezurawska04}, demonstrating that 1720\\,MHz masers may occur under different pumping conditions and OH column densities than those present in SNRs \\citep{gray91,gray92}. Concentrating on the SNRs, the masers originate in the shocked region where the expanding SNR collides with the surrounding interstellar medium (ISM). Consistent with this model, a large number of 1720\\,MHz masers are found in the Galactic center (GC) where the Sgr\\,A\\,East SNR plows into the ISM, notably into the $+50$~\\kms\\ molecular cloud \\citep[e.g.,][hereafter Paper\\,I]{sjouwerman08}. The majority of these masers are observed along a circular pattern outlined by the expansion of the SNR, displaying a relatively small range of line of sight velocities between $+34\\leq V_{\\rm LSR}\\leq 66$\\,km\\,s$^{-1}$ \\citep[e.g.,][]{yusefzadeh96,karlsson03,sjouwerman08}. In addition, two separate groups of masers are located near the Circumnuclear Disk (CND) that are not directly explained by the SNR/ISM interaction model. These two groups of masers have velocities that are offset from the Sgr\\,A\\,East masers; $V_{\\rm LSR}\\simeq+130$\\,km\\,s$^{-1}$ and $V_{\\rm LSR}\\simeq-130$\\,km\\,s$^{-1}$ respectively (Paper\\,I). In Paper\\,I we argue that these masers are unlikely to be pumped by a shock produced by Sgr\\,A\\,East. Other plausible pumping scenarios include local shocks produced by random motions of clumps or turbulence, supported by the presence of strong H$_2$ (1-0) $S(1)$ emission \\citep{yusefzadeh01}, or by infrared (IR) pumping similar to conditions observed in star forming regions \\citep{gray92}.   The environment in the CND is likely to be different from that in a SNR/ISM post-shock region. Variability studies of masers can be used to probe the environment and could further shed light on the differences in excitation mechanisms and conditions for CND versus SNR/ISM masers. In the CND, the pumping may include IR pumping routes, since local IR peaks are observed within the CND \\citep{latvakoski99}. This would be in contrast to the 'standard' SNR masers, which are pumped by collisions only. Also, differences in collision rate along the line of sight, due to clumpiness of the medium, could result in a different maser flux variability of the CND masers as compared to the SNR/ISM masers \\citep{pihlstrom08,pihlstrom01}.   \\begin{deluxetable}{rlccc}[th] \\tabletypesize{\\scriptsize} \\tablecaption{Observational details of the VLA monitoring program performed in 2006} \\tablewidth{0pt}  \\tablehead{ \\colhead{$\\phantom{0000}$Day of} & \\colhead{Config.} & \\colhead{Beam size} & \\colhead{Channel rms} & \\colhead{Obs.\\ time}\\\\ \\colhead{$\\phantom{0000}$Year} & & \\colhead{\\arcsec$\\times$\\arcsec} & \\colhead{mJy\\,beam$^{-1}$} & \\colhead{hours} } \\startdata 1$\\phantom{0000}$47 & A   & 2.7$\\times$1.2 & 5.2 & 5 \\\\ 2$\\phantom{000}$144 & BnA & 3.9$\\times$2.8 & 6.3 & 4 \\\\ 3$\\phantom{000}$164 & BnA & 5.7$\\times$3.3 & 7.2 & 4 \\\\ 4$\\phantom{000}$196 & B   & 7.7$\\times$3.8 & 8.0 & 4 \\\\ 5$\\phantom{000}$242 & B   & 7.6$\\times$3.4 & 6.6 & 4 \\\\ \\enddata \\label{observations} \\end{deluxetable}    Variability of other OH maser transitions has been observed. For example, the 1612, 1665 and 1667 MHz circumstellar OH masers in variable stars frequently exhibit variability that is correlated with the stellar cycle \\citep[e.g.,][]{jewell79, etoka00}. Thus, this variability can be directly coupled to the infrared pumping mechanism of 1612, 1665 and 1667 MHz masers. For 1720 MHz masers in star forming regions, variability over two epochs have been reported by \\citet{caswell04}, and variability on 15-20 minute timescale is seen for the masers in the star forming region W3(OH) \\citep{ramachandran06}. However, the pumping of 1720 MHz masers in star forming regions include radiative pumping routes via IR emission from the central star \\citep{gray91, gray92}. Such radiative pumping scheme is not available in the SNRs, which are completely collisionally pumped \\citep{elitzur76, lockett99, wardle99, pihlstrom08}. Since these masers are excited by the SNR shock, it is less likely that any observed variability would be induced by pumping variations. Instead we expect that variations in the amplification path length would be the most direct method to produce amplitude variations in SNR/ISM 1720 MHz masers.  As little is known about the 1720 MHz maser variability in SNRs, we here present a five-epoch variability study of the 1720 MHz masers in the GC with the Very Large Array (VLA). ",
        "conclusions": "\\label{discussion} During the 195 days over which our study took place the most significant result is that the northeastern masers associated with the SNR/ISM interaction show a higher degree of variability than the rest of the masers. Over the observed timescale no masers disappeared, and no new masers appeared. Compared to our observations in 2005 (Paper\\,I) one new maser appeared (or more precisely, this time fulfilled our detection criteria). The two ``archival'' masers, included for completeness in Paper\\,I, remain undetected. This indicates a quite stable environment for both SNR/ISM and CND masers during our monitoring campaign. In this section we first discuss how variability can be introduced by both external and intrinsic (Sect.\\ \\ref{intr}) sources. Thereafter, we explore the nature of the variability of the SNR (Sect.\\ \\ref{snrvar}) and CND (Sect.\\ \\ref{cndvar}) masers respectively.   \\subsection{Maser variability} \\label{intr}   Observational errors such as a systematic incorrect value in the derived value of our flux calibrator in one of our epochs, pointing errors, etc., will have a predictable signature in the individual measurements. We have carefully looked for these in both our target and calibrator data and conclude that none of them contribute to the overall results. Flux variations can also be produced by interstellar scintillation \\citep{narayan92,cordes98}. This causes variations on timescales of a few minutes. Such short timescales are not considered in our study since in data averaged over a minute only a couple of masers would be bright enough to be reliably detected.  Assuming that the variability of the maser intensity is not due to observational effects, variability studies can be used to probe the maser environment. Changes in maser peak velocity could indicate acceleration, such as observed in outflows \\citep{liljestrom89, brand03}. In stellar masers, periodic maser variations may be related to periodicity found in the central object, in turn affecting the pumping conditions and/or the maser path length. SiO masers in the evolved Mira variables are examples of how the maser luminosity is a result of the phase of the central star \\citep[e.g.][]{wittkowski07}. Another example can be found in star forming regions (SFRs), where 22~GHz water masers are well-known to exhibit maser variability. The variability of these masers are characterized by large (sometimes several orders of magnitude) amplitude variations, and often with velocity drifts of a few \\kms\\ per year \\citep{comoretto90, wouterloot95, claussen96, brand03, furuya01, furuya03}. This occurs on timescales from a few hours to years. Water masers associated with SFRs are excited behind shocks, presumably caused by outflows or jets driven by the young stellar object (YSO), and much of the variation can be attributed to changes in the luminosity of the YSO, in turn affecting the pumping rate via changes in the outflows and jets \\citep{elitzur89, felli92}.  OH masers associated with SFRs are predominantly observed in the radiatively pumped 1665 and 1667 MHz lines (sometimes accompanied by masers in the 1612 and 1720\\,MHz transitions), and variability of those masers have been investigated by multiple groups \\citep{schwartz74, zuckerman72, rickard75, gruber76, clegg91}. The long term variability (weeks to months) can be attributed to changes in the number density of inverted OH molecules, or path length changes. In contrast, the short term variability is assumed to be related to sudden changes in the pumping mechanism, reflecting fluctuations in the host star luminosity. For OH masers in SNRs, the situation has been unknown. The 1720~MHz OH transition was for a long time the only transition observed near SNRs \\citep{lockett99, wardle99, pihlstrom08}, although recent observations have shown Sgr\\,A to harbor 36.2 and 44.1~GHz methanol masers \\citep{sjouwerman10,pihlstrom11}. The 1720~MHz OH masers are usually assumed to originate in the post-shock region where the expanding SNR collides with the surrounding medium. No dedicated variability studies of 1720~MHz OH masers have been published to date, and we therefore know little about the stability of the maser environment in these sources. Here we have a shocked environment and since this is a large scale phenomena it must arise in the maser generating column density, or from the subsequent propagation of emission through the interstellar medium.    \\subsection{The SNR masers} \\label{snrvar}   The northeastern SNR masers clearly are variable on the timescales of weeks to months. These masers are the result of the SNR Sgr\\,A\\,East ramming into the $+50$~\\kms\\ cloud and we expect the post-shock medium to display density variations due to turbulence in those regions, given that a shock has recently passed. Assuming that the masers can be considered as a series of velocity coherent regions of gas, bulk gas flow and turbulent motion in the source will move parts of the maser column in and out of velocity coherence along the line of sight. As a result, the maser will exhibit intensity variations. For the northeastern masers this cannot be the only source of the variability though, since much higher levels of variability are observed than toward other regions of Sgr\\,A\\,East. It is not clear whether it is the maser generating column density itself or the medium in front of the masers that is causing the variability. However, the $+50$~\\kms\\ cloud is thought to be located behind Sgr\\,A\\,East \\citep[e.g.,][Paper\\,I]{coil00} and thus between the observer and the masers. When observing this group of masers we therefore probably look through a longer path length of turbulent gas than toward other regions of Sgr\\,A\\,East. This is supported by OH absorption observations showing this region harbors a much higher absorbing OH column density than the southern and northwestern interaction regions (Paper\\,I, Karlsson et al.\\ 2003). It is possible that the supernova covering the northeastern region contributes to the turbulence of this column density.  In contrast, the southern masers are occurring in a region where the two SNR Sgr\\,A\\,East and G359.02$-$0.09 are colliding. Here there is no clear evidence of large columns of gas in front of the masers, and no supernova crossing the line of sight.    \\subsection{The CND masers} \\label{cndvar}  The CND masers exhibit much less variability than the northeastern SNR masers, with no maser displaying variability with a $VI>1$. If the northeastern SNR/ISM maser variability is due to passage of the maser through a screen of hot turbulent gas, the (non-)variability of the CND masers is consistent with CND located {\\em in front of} Sgr\\,A\\,East (Paper\\,I).  \\begin{figure}[th] \\resizebox{8.5cm}{!}{\\rotatebox{0}{\\includegraphics{f2.eps}}} \\caption{Position of each maser listed in Table \\ref{fluxes}, in relation to the 1.7\\,GHz continuum.} \\label{maserpos} \\end{figure}  Since variability could give insight into the maser environment, we consider whether the low variability is consistent with the presumed pumping and amplification conditions in the CND. Here the pumping cannot be due to a SNR/ISM shock, but the CND masers are thought to be pumped by collisions between clumps in the CND (Paper\\,I). The number of clumps within the line of sight are probably only a few since once a clump has moved out of the line of sight, a maser is likely to disappear completely. Assuming an inclination of 65$^\\circ$ \\citep{latvakoski99,marr93,jackson93,christopher05} and thus a rotational velocity of approximately 160~\\kms\\ of the maser emitting clouds in the CND, the motion in the plane of the sky over the 195 days is minimal; only $3.7\\times 10^{-5}$ pc (approximately 8 AU). This is much smaller than the high density cloud core sizes of $\\sim 0.25$ pc that are estimated by molecular line observations \\citep{christopher05}.  Similarly, assuming typical relative clump velocities of 20-30~\\kms\\ in the sky-plane and 0.1~pc in size, a clump will pass by a compact maser in $\\sim$3900 years. For a clump size of 1~AU the corresponding time is 2 months. If clump collisions and motions are involved in determining the path lengths for the CND masers, given the low variability levels observed the clumps must be much larger than 1~AU in size. More likely they are only a fraction of the size of the dense molecular clumps observed in HCN, perhaps corresponding to a high density region of the clump. For such clump sizes little variability can be expected due to path length changes.  A factor that could contribute to a more stable maser gain in the CND region is the pumping conditions. As measured in molecular line emission, HCN and HCO$^+$ line intensity is stronger in the CND than in regions associated with, for example, the $+50$~\\kms\\ cloud \\citep{wright01}, implying higher densities. A higher density would result in a higher collision rate, thus keeping the pumping at a higher rate eventually resulting in more stable maser intensities. Note that \\cite{yusefzadeh01} show that densities $n>10^6$ cm$^{-3}$ are needed for clump-clump collisions to produce 1720\\,MHz pumping, which is slightly higher than the C-shock post-region density of $n\\simeq10^5$ cm$^{-3}$ for the SNR masers \\citep{lockett99,wardle99}.   What the exact pumping conditions and the pumping rate really are is difficult to calculate, and depends on the ortho-para $H_2$ ratio in the region. Both \\citet{lockett99} and \\citet{pavlakis96} have found that collisions with para-H$_2$ can strongly suppress the 1720 MHz inversion. Ortho- and para-H$_2$ are thought to be formed on grains with a ratio of 3:1.  At low temperatures ($T=10$ K) proton exchange reaction convert ortho-H$_2$ into para-H$_2$, making the para-H$_2$ being the dominant species. As the temperature increases, there is less of a difference, but still with para-H$_2$ dominating \\citep{offer92}. However, to create the 1720~MHz maser only a small amount of ortho-H$_2$ is required, since the collision rate for ortho-H$_2$ usually is larger by a factor of 2-3 than the para rates \\citep{offer92}. The equilibrium ratio for ortho-para H$_2$ is $9\\times e^{-170/T}$, which is equal to 1 for a temperature $T=77$\\,K, implying the temperature should be cooler than this value. This is consistent with a high density environment in the CND where the gas temperature is thermalized with the dust temperatures of 20-80\\,K \\citep{becklin82,mezger89}, although molecular studies of the CND imply excitation temperatures anywhere between 50-200\\,K \\citep{christopher05,jackson93,wright01,coil99,coil00}. \\citet{coil00} use NH$_3$(2,2)/NH$_3$)(1,1) ratios to derive rotational gas temperatures between 20-70~K, mostly in the outer regions of the CND. Perhaps, next to path-length geometry (Paper\\,I), this is a reason that 1720 MHz masers are not seen all across the CND; if temperatures are too high the 1720 MHz emission will be suppressed."
    },
    "1107/1107.4177_arXiv.txt": {
        "abstract": "We surveyed the (1,1), (2,2), and (3,3) lines of NH$_3$ and the H$_2$O maser toward the molecular cloud L935 in the extended H\\emissiontype{II} region NGC 7000 with an angular resolution of \\timeform{1.6'} using the Kashima 34-m telescope. We found five clumps in the NH$_3$ emission with a size of 0.2--1 pc and mass of 9--452 \\MO. The molecular gas in these clumps has a similar gas kinetic temperature of 11--15 K and a line width of 1--2 km s$^{-1}$. However, they have different star formation activities such as the concentration of T-Tauri type stars and the association of H$_2$O maser sources. We found that these star formation activities are related to the geometry of the H\\emissiontype{II} region. The clump associated with the T-Tauri type star cluster has a high star formation efficiency of 36--62\\%. This clump is located near the boundary of the H\\emissiontype{II} region and molecular cloud. Therefore, we suggest that the star formation efficiency increases because of the triggered star formation. ",
        "introduction": "The star formation efficiencies (SFEs) of molecular clouds in the Milky Way Galaxy, are typically observed to be 10\\%. The SFEs in nearby molecular clouds are $\\simeq$ 3--15\\% (\\cite{swift2008}; \\cite{evans2009}). The observations of giant molecular clouds in the inner Galaxy indicate that the SFEs in these clouds are of the order of a few percent \\citep{myers1986}. However, \\citet{lada1992} found that three of five massive cores, NGC 2024, NGC 2068, and NGC 2071 exhibit higher SFEs of $\\simeq$ 30--40\\%. It remains unclear why these SFEs are high. \\citet{lada1992} suggested that the high gas densities and high gas mass may be required for the high SFE but there should be additional conditions, because the other two cores exhibit low SFEs of $\\simeq 7$\\%. However, additional conditions for high SFE are unknown, and more observational investigations are required. \\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnote[1]{Present address: Graduate School of Science and Engineering, Kagoshima University, 1-21-35 Korimoto, Kagoshima, Kagoshima 890-0065.}  NGC 7000 is an extended H\\emissiontype{II} region in the Cygnus X region. On its southeastern side is a molecular cloud L935. The $^{13}$CO emission in this molecular cloud is the brightest in the Cygnus X region \\citep{dobashi1994}. Figure \\ref{fig:1} shows the optical image of NGC 7000 and L935. Seven T-Tauri stars are clearly clustered at the boundary of the H\\emissiontype{II} region \\citep{herbig1958}. This suggests that the star formation is triggered by the interaction of the H\\emissiontype{II} region with the dense molecular gas. A number of studies of star formation in a cloud associated with an H\\emissiontype{II} region have been performed (\\cite{sugitani1989}; \\cite{sugitani1991}; \\cite{sugitani1994}; \\cite{dobashi2001} \\cite{deharveng2003}; \\cite{deharveng2005}). For example, in the nearby H\\emissiontype{II} region, IC 5070, a molecular shell with an expanding velocity of 5 km s$^{-1}$, is found in the $^{12}$CO ($J$=1-0) line \\citep{bally1980}. They suggest that the T-Tauri type stars in IC 5070 are formed by the expanding shell.  Our aim is to investigate the relationship between dense molecular gas and star formation based on the SFE. NGC700 has the advantage of allowing us to estimate the SFE because T-Tauri type stars are associated with it and we can estimate the stellar mass accurately. Therefore, we made the observations in the NH$_3$ line to estimate the mass of dense molecular gas. We also surveyed an H$_2$O maser source that is associated with outflow from a young stellar object (YSO). We adopted the distance to NGC 7000 to be 600 pc \\citep{laugalys2002}. ",
        "conclusions": "We observed NGC 7000 in the NH$_3$ line and H$_2$O maser using the Kashima 34-m telescope. Our observations are summarized as follows:  \\begin{enumerate}  \\item We found two major clumps with a mass of 95--452 \\MO, and three subclumps with a mass of 9--20 \\MO. The molecular gas in these show similar gas kinetic temperatures of 11--15 K and line width of 1--2 km s$^{-1}$. However, they show different star formation activities such as the concentration of T-Tauri type stars and the association of an H$_2$O maser.  \\item One of the clumps that is associated with a cluster of T-Tauri type stars shows the SFE $\\simeq$ 36--62\\%. This SFE is higher than that of the other clumps.  \\item A comparison of the distribution of molecular gas and ionized gas traced by the H$\\alpha$ emission suggests that the clump with high SFE is located near the H\\emissiontype{II} region. Therefore, the high SFE would be related to the interaction of molecular gas and the H\\emissiontype{II} region.  \\item We found a new H$_2$O maser source in the NH$_3$ clump. Although the counterpart of this maser is not found in the IRAS point source catalogue, we found it in the $Spitzer$ 24- and 70-$\\mu$m images. This suggests that a new H$_2$O maser survey  should be carried out based on the point source catalogue of $Spitzer$ and/or AKARI.  \\end{enumerate}   \\bigskip  We thank an anonymous referee for very useful comments and suggestions. T.O. was supported by a Grant-in-Aid for Scientific Research from the Japan Society for the Promotion Science (17340055). We acknowledge K. Miyazawa (NAOJ) for his technical support of observations."
    },
    "1107/1107.3561_arXiv.txt": {
        "abstract": "Sub-mm observations of protoplanetary disks now approach the acuity needed to measure the turbulent broadening of molecular lines. These measurements constrain disk angular momentum transport, and furnish evidence of the turbulent environment within which planetesimal formation takes place. We use local magnetohydrodynamic (MHD) simulations of the magnetorotational instability (MRI) to predict the distribution of turbulent velocities in low mass protoplanetary disks, as a function of radius and height above the mid-plane. We model both ideal MHD disks, and disks in which Ohmic dissipation results in a dead zone of suppressed turbulence near the mid-plane. Under ideal conditions, the disk mid-plane is characterized by a velocity distribution that peaks near $v \\simeq 0.1 c_s$ (where $c_s$ is the local sound speed), while supersonic velocities are reached at $z > 3 H$ (where $H$ is the pressure scale height). Residual velocities of $v \\approx 10^{-2} c_s$ persist near the mid-plane in dead zones, while the surface layers remain active. Anisotropic variation of the linewidth with disk inclination is modest. We compare our MHD results to hydrodynamic simulations in which large-scale forcing is used to initiate similar turbulent velocities. We show that the qualitative trend of increasing $v$ with height, seen in the MHD case, persists for forced turbulence and is likely a generic property of disk turbulence. Percent level determinations of $v$ at different heights within the disk, or spatially resolved observations that probe the inner disk containing the dead zone region, are therefore needed to test whether the MRI is responsible for protoplanetary disk turbulence. ",
        "introduction": "Understanding the structure of protoplanetary disks is central to modeling the phenomenology of Young Stellar Objects \\citep{williams11} and to theoretical studies of all phases of the formation of planetary systems. For a significant fraction of their lives, gas within protoplanetary disks is observed to actively accrete onto the central star, probably as a consequence of turbulent transport of angular momentum \\citep{shakura73}. Globally, this accretion and redistribution of angular momentum results in evolution of the surface density profile \\citep{lyndenbell74}, which limits the time scale for massive planet formation and affects quantities such as the rate of planet migration \\citep{lubow10}. On smaller scales, turbulence sets the environment for planetesimal formation \\citep{chiang10} by determining both the local concentration and collision velocities \\citep{volk80,ormel07} of small particles that are aerodynamically coupled to the gas.  Although several physical processes --- including self-gravity and the magnetorotational instability \\cite[MRI;][]{balbus98} --- may initiate disk turbulence \\citep[for a review, see][]{armitage11}, existing theoretical studies have largely been untroubled by observational validation. The most widely accepted constraint on disk turbulence comes from measurements of disk lifetimes \\citep{haisch01} and accretion rates \\citep{hartmann98}, which imply that protoplanetary disks around low-mass stars evolve and are dispersed on Myr time scales. This observation pins down the angular momentum transport efficiency if the evolution results from turbulence; the efficiency is conventionally expressed in terms of a \\cite{shakura73} $\\alpha \\approx 10^{-2}$. Generically, this level of stress within a fluid disk implies characteristic velocity perturbations $v \\sim \\alpha^{1/2} c_s \\sim 0.1 c_s$ \\citep[where $c_s$ is the sound speed, e.g.][]{balbus98}, but this estimate is so crude as to be useful mainly for motivating further observations. Neither it, nor other constraints on $\\alpha$ from detailed modeling of individual systems \\citep{hueso05} provide any information on the nature of turbulence or on any dependence of its properties on height above the mid-plane.  Direct determination of the strength of protoplanetary disk turbulence is possible by detecting the turbulent broadening of molecular lines observed in the infrared \\citep{carr04} or sub-mm \\citep{hughes11}. Subsonic turbulent broadening is a challenging quantity to measure, as protoplanetary disks are comprised of supersonically orbiting gas; thus, precise measurements are needed to separate the small turbulent component from the dominant bulk rotation. Furthermore, in the inner disk, observed lines from the disk may be contaminated by outflow components \\citep{bast11}. Nonetheless, current observations of the outer regions of disks already attain precisions comparable to the level ($v \\sim 0.1 c_s$) where a signal can plausibly be expected. Using the Submillimeter Array ({\\em SMA}) to observe the CO(3-2) transition, \\citet{hughes11} derived constraints on the turbulent linewidth in the atmosphere of the disks surrounding the T~Tauri star TW~Hya and the Herbig~Ae star HD~163296\\footnote{These are not ``typical\" sources. TW~Hya is a nearby system with a favorable near face-on geometry, while HD~163296 has a very large (500~AU) disk.}. For TW~Hya they placed an upper limit to the turbulent velocity of $v < 0.1 c_s$, while for HD~163296 they obtained a tentative detection of turbulent broadening corresponding to $v \\approx 0.4 c_s$. Although still preliminary, these observations provide a clear indication that {\\em ALMA}, with superior sensitivity and spatial resolution, will constrain disk turbulence to theoretically interesting levels for these and other sources.  In this paper, our goal is to quantify the expected turbulent velocities in protoplanetary disks as a function of radius and height above the mid-plane. We focus on low mass disks (TW~Hya would be a good example) which are stable against self-gravity, and assume that the MRI is the sole source of turbulence. We compute both reference models in the ideal magnetohydrodynamic (MHD) limit, and physical models in which the MRI is partially damped by Ohmic dissipation, forming a dead zone \\citep{gammie96,sano00,fromang02}. Particular care is taken to ensure that the results are numerically converged; we use local shearing box simulations whose convergence with spatial resolution has previously been demonstrated \\citep{davis10,simon11}, and we explicitly check the effect of varying the domain size. We also calculate purely hydrodynamic simulations in which turbulence is initiated through arbitrary large-scale forcing. By comparing these to the MHD runs, we address the question of whether the observable properties of disk turbulence can constrain the underlying mechanism that initiates angular momentum transport.  The plan of the paper is as follows. In \\S2 we describe the numerical simulations in ideal MHD, non-ideal MHD, and pure hydrodynamics that form the basis of the turbulent velocity calculation. In \\S3 we outline the velocity distribution calculation, the results of which are shown in \\S4. \\S5 discusses our results and their implications for future observations. Finally, we summarize our conclusions in \\S 6. ",
        "conclusions": "\\label{conclusions}  Our conclusions about the turbulent properties of low mass protoplanetary disks are: \\begin{enumerate}  \\item Characteristic turbulent velocities are $\\sim$ (0.1-1)$\\cs$ for fully turbulent regions, in rough agreement with observations made to date \\cite[e.g.,][]{hughes11}. These characteristic velocities are reasonably robust to variations in numerical (changes to local domain size) and physical (locations in a model disk) parameters.  \\item Turbulent velocity increases away from the disk mid-plane due to steepening.  In the upper region of the disk ($|z| \\gtrsim 3 H$), the velocity distribution peaks around $0.5\\cs$ and has a significant ($\\sim 10\\%$) supersonic component.  As one probes towards the mid-plane, $|v|/\\cs \\sim 0.1$ is typical of fully turbulent disks.  \\item In calculations with an MRI dead zone near the mid-plane, the characteristic turbulent velocities are $\\sim 0.01\\cs$ within the dead zone.  In principle, with an improvement in sensitivity, observations that probe different depths could see the presence of the dead zone as velocities drop from $\\sim$ 0.1-1$\\cs$ to $\\sim 0.01\\cs$.  \\item The density structure for $|z| > 2H$ is significantly different in the MRI versus purely hydrodynamic cases, which could have potential implications for the observed turbulent linewidths if different depths can be probed.  \\item The vertical and planar velocity distributions are quite similar, suggesting that turbulent linewidths will only weakly be dependent on the inclination angle.  \\end{enumerate}  Our predictions for the distribution of turbulent velocity in MRI-active disks suffer from a number of uncertainties. First, we have chosen to only focus on one non-ideal MHD effect, namely Ohmic resistivity.  Other non-ideal effects -- ambipolar diffusion and the Hall term -- are also important in protoplanetary disks \\citep{kunz04,bai11,wardle11}. Ambipolar diffusion, in particular, is important in low density regions, and may affect the properties of turbulence in the most observationally accessible location -- the disk atmosphere at large radius. Moreover, the resistivity that we have employed neglects dust physics. We also caution that some of the most striking qualitative trends that we observe are linked to the steepening of waves near the disk surface. Wave propagation in disks is known to depend upon the vertical thermal structure \\citep{bate02}, and hence the isothermal structure that we have assumed may not always be adequate. Even before adding a treatment of the radiation physics, these limitations imply that there remains much work to be done to further constrain the turbulent velocities in simulations of MRI-active disks.  While our results coupled with recent observations provide mild support for the model of MRI-driven angular momentum transport, the calculations that we have presented here have not been able to identify a strong observational discriminant between MRI-driven and purely hydrodynamic turbulence. Quite precise measurements of the turbulence as a function of height will be needed to tell one from the other on purely observational grounds.  In fact, if sufficiently high spatial resolution observations of the inner disk regions reveal the presence of a dead zone region, then this would present very strong support for the MRI driving disk turbulence.  Finally, we reiterate that we have considered only arbitrary hydrodynamic forcing, rather than setting up known physical drivers of turbulence, such as self-gravity or even convection \\citep{lesur10}.   By design, the average kinetic energy in the purely hydrodynamic simulations nearly equals that of the MRI simulations. If it could be established, theoretically, that hydrodynamic drivers of turbulence were necessarily weaker than the MRI, a single measurement of the turbulent velocity would then distinguish between the two. If, on the other hand, we treat the strength of hydrodynamic turbulence as a free parameter, then our results suggest that the vertical variation of turbulent velocities in the hydrodynamic and MHD limits can have a qualitatively similar trend. Of course, the physical mechanisms that might initiate hydrodynamic turbulence without arbitrary forcing could, in principle, imprint distinctive characteristics into the observable turbulent velocity field, which would allow them to be distinguished from the MRI more readily. To test this, it would be useful to repeat the analysis presented here for disks in which these other sources of turbulence are active. These calculations are currently underway and will be presented in future work."
    },
    "1107/1107.4207_arXiv.txt": {
        "abstract": "One key feature of the interacting stellar winds model of the formation of planetary nebulae (PNe) is the presence of shock-heated stellar wind confined in the central cavities of PNe. This so-called hot bubble should be detectable in X-rays. Here we present \\emph{XMM-Newton} observations of NGC\\,3242, a multiple-shell PN whose shell morphology is consistent with the interacting stellar winds model. Diffuse X-ray emission is detected within its inner shell with a plasma temperature $\\sim$2.35$\\times$10$^6$~K and an intrinsic X-ray luminosity $\\sim$2$\\times$10$^{30}$ ergs~s$^{-1}$ at the adopted distance of 0.55 kpc. The observed X-ray temperature and luminosity are in agreement with ``ad-hoc'' predictions of models including heat conduction. However, the chemical abundances of the X-ray-emitting plasma seem to imply little evaporation of cold material into the hot bubble, whereas the thermal pressure of the hot gas is unlikely to drive the nebular expansion as it is lower than that of the inner shell rim. These inconsistencies are compounded by the apparent large filling factor of the hot gas within the central cavity of NGC\\,3242. ",
        "introduction": "Planetary Nebulae (PNe) consist of stellar material ejected by low- and intermediate-mass stars (0.8--1.0~M$_\\odot\\le M_i \\le$\\,8--10~M$_\\odot$). Towards the end of the Asymptotic Giant Branch (AGB), these stars experience copious mass loss and eject most of their stellar envelope through a slow, dense AGB wind. The ejected material is subsequently ionized by the central star and becomes a PN. PNe eventually disperse into the interstellar medium as they expand, whereas the stellar cores, mainly composed of carbon and oxygen, will evolve toward the white dwarf stage.   Near the time when the hot stellar core is exposed, the slow AGB wind, with terminal velocities 5--30 km s$^{-1}$ \\citep{ELT88}, is superseded by a fast stellar wind with terminal velocities 1,000--4,000 km s$^{-1}$ \\citep{CRP85,GR-LM10}. This fast stellar wind sweeps up the slower AGB wind to form a PN \\citep{Kwok83}. In this interacting stellar winds (ISW) model, the physical structure of a PN would be similar to that of a wind-blown bubble, as modeled by \\citet{WR77}, comprising a central cavity filled with shocked fast wind (the so-called hot bubble), a dense shell of swept-up AGB wind at 10$^{4}$ K, and an outer envelope of unperturbed expanding AGB wind. In a simplistic model, the temperature of the shocked stellar wind inside the hot bubble would be 10$^{7}$-10$^{8}$ K, but turbulent mixing \\citep[e.g.,][]{MF95} or heat conduction \\citep{ZP98,SSW08} lowers the temperature of the hot gas to 10$^{6}$-10$^{7}$ K and raises its density to produce optimal conditions for the emission of soft X-rays. Therefore, X-ray observations of shock-heated hot gas in PNe provide us a direct means to examine the interaction of the fast stellar wind with the AGB wind and to investigate the transfer of energy and momentum to the PN envelope.   \\emph{ROSAT} observations showed hints of diffuse X-ray emission in a few PNe \\citep{GR00}. However, it was not until the advent of \\emph{Chandra} and \\emph{XMM-Newton}, with their unprecedented resolution and sensitivity, that we were finally able to unambiguously detect hot gas in PNe. \\emph{Chandra} and \\emph{XMM-Newton} have resolved the diffuse X-ray emission in a handful of PNe \\citep[e.g.,][]{KN00,KN01,CH01,GR02,GR05} and revealed unexpected, hard X-ray emission from the central stars of several PNe that may originate from the coronal emission of unseen faint binary companions or shocks within the fast stellar winds \\citep{GR01,KN03,Montez_etal2010}. Observations of diffuse X-ray emission in PNe demonstrate that hot gas in elliptical PNe is confined within the innermost nebular shell and that the high pressure of the hot gas may indeed drive the nebular expansion as expected in bubble models.   The observed values of $L_{\\rm X}$ and $T_{\\rm X}$ are in agreement with predictions of the time-dependent models developed by \\citet{SSW08} that include heat-conduction for PNe with central stars of normal, hydrogen-rich surface composition. There is, however, an unsolved puzzle: the analyses of the chemical composition of the X-ray-emitting plasma in PNe suggest that it is mainly composed of stellar wind, with little contamination of material from the cold nebular shell \\citep[e.g., NGC\\,2392 and NGC\\,6543,][]{GR05,CH01}. This discrepancy has been further illustrated by the analysis of the high-spectral resolution \\emph{Chandra} LETG spectrum of BD+30\\degr3639 that conclusively confirms the stellar wind composition of its X-ray-emitting plasma \\citep{Yu_etal2009}. This problem has prompted alternative mechanisms for the production of hot gas in PNe, including the action of fast collimated outflows and/or slow fragments in the onset of the fast stellar wind \\citep{SK03}, and the absorption of energy from the stellar wind by slowly moving ions embedded in the wind itself, the so-called pick-up ions \\citep{Soker_etal2010}.    NGC\\,3242 (PN\\,G261.0+32.0), the Ghost of Jupiter, is a multiple-shell PN with a bright, 28\\arcsec$\\times$20\\arcsec\\ inner ellipsoidal shell and ansae surrounded by a fainter, 46\\arcsec$\\times$40\\arcsec\\ moderately elliptical envelope. These two shells are further enclosed by arcs and a giant broken halo revealed by deep images \\citep{COR03,COR04}. The double-shell morphology of the main nebula of NGC\\,3242 is highly suggestive of interactions between the fast stellar wind of its central star \\citep[$v_{\\infty}$=2,400 km~s$^{-1}$,][]{PD04} and the previous slow AGB wind. The AGB wind (the nebular envelope) has been swept by the fast stellar wind to form a thin ionized shell with a central cavity that can be expected to be filled with shocked fast wind. This shock-heated gas should emit X-rays, and the diffuse X-ray emission from NGC\\,3242 is likely detectable because of its proximity \\citep[distance = 0.55$\\pm$0.23 kpc,][]{TZ97,ML04}, and low extinction \\citep{Balick_etal93,HE00,POT08}.   In this paper, we present \\emph{XMM-Newton} observations of NGC\\,3242 that have detected diffuse X-ray emission within its innermost nebular shell. The observations are described in \\S2, the results are presented in \\S3, the physical structure of the optical nebula is investigated in \\S4, and the effects of the shocked stellar wind in the nebula are discussed in \\S5. ",
        "conclusions": "The properties of the X-ray emission from NGC\\,3242, based on a preliminary analysis \\citep{Ruiz_etal06}, were compared to 1D hydrodynamical models of PNe that included the effects of stellar wind evolution and heat conduction by \\citet{SSW08}. They remarked that the round shape of the rim of the inner shell of NGC\\,3242 makes it a suitable PN for the comparison with their 1D simulations. The revised values of different X-ray properties of NGC\\,3242 presented here need to be discussed in the framework of \\citet{SSW08}'s model. This revision mostly affects the values of the thermal pressure of the hot bubble and inner shell rim, whereas the X-ray luminosity (scaled to the distance of 0.55 kpc) and plasma temperature of NGC\\,3242 are basically the same as those used by \\citet{SSW08}.   We note that the models selected by \\citet{SSW08} have a thermal pressure of the hot bubble that exceeds that of the rim, while our revised values indicate the opposite, i.e., the thermal pressure of the rim ($P_{\\rm th}$=3.6$\\times$10$^{-9}$ dyne~cm$^{-2}$) is higher than that of the bubble ($P_{\\rm th}$=1.6$\\times$10$^{-9}\\,\\epsilon^{-1/2}\\,d^{1/2}$ dyne~cm$^{-2}$). One possible outcome to this issue would be a distance for NGC\\,3242 much farther than 0.55 kpc \\citep{TZ97,ML04}; at a distance of 3 kpc, the thermal pressures of the hot bubble and rim would be the same, but this would make NGC\\,3242 too large ($r=0.12$ pc) and evolved ($\\tau=6,000$ yrs) and we consider it to be unlikely.   The other possible solution is to consider a low value for the filling factor of the X-ray-emitting gas in the bubble; a value $\\epsilon$=0.15 would make both pressures the same. Such a low filling factor is expected to produce a noticeable limb-brightening morphology that is not observed (Figures~\\ref{ximg} and \\ref{prof}). As shown in Fig.~\\ref{sim_ximg}, hints of a limb-brightened morphology are still apparent for a hot gas shell with a fractional thickness of 0.2 (i.e., $\\epsilon \\sim 0.5$) even at the \\emph{XMM-Newton} limited spatial resolution ($\\sim15^{\\prime\\prime}$). It is arguable, however, that extinction has reduced the center-to-limb contrast of the X-ray emission, thus shifting the emission inwards, as this is an effect expected to be noticeable \\citep{SSW08} even for the small extinction ($N_{\\rm H}$=5$\\times$10$^{20}$ cm$^{-2}$) towards the nebula.   Alternatively, we must consider the dynamical effects of the photo-ionization of the nebular material, which may have overcome those of the currently diminished stellar wind of NGC\\,3242 \\citep{Kudritzki_etal97,TL02}. In this case, the nebular expansion is driven by the thermal pressure increase produced in the rim by photo-ionization. This may explain the relatively large thickness of NGC\\,3242 inner shell rim, as compared to that of NGC\\,6543. We suggest that the rim thickness can be used as an indicator of the relative thermal pressure of hot gas and photo-ionized nebular material; a sharp rim would imply that the thermal pressure of the hot gas is larger than this of the nebula, and thus that hot gas itself is present (and detectable), while a thick inner shell rim would imply a lower pressure for the hot gas, if present.   The low plasma temperature is puzzling in view of the enhanced N/O ratio of the X-ray-emitting plasma, about three times higher than the N/O ratio of the nebular material. Such differences in the chemical abundances seem to imply that the X-ray-emitting plasma mostly consists of shocked stellar wind, with little contamination of material from the nebula. The low plasma temperature, however, needs large amounts of material from the cold nebular shell to have been incorporated into the hot bubble. This problem, similar to that found in other PNe such as BD+30\\degr3639 \\citep{Yu_etal2009}, NGC\\,2392 \\citep{GR05}, and NGC\\,6543 \\citep{CH01}, may require mechanisms for the production of X-ray emission in PNe other than mixing and heat conduction as summarized by \\citet{Soker_etal2010}.   There is, finally, a noticeable asymmetry of the spatial distribution of the X-ray emission from NGC\\,3242, with the northwestern half of the hot bubble being brighter than its southeastern half (Figures~\\ref{ximg} and \\ref{prof}). Even if we assume that the inner shell of NGC\\,3242 were a tilted ellipsoid with its Northwestern tip moving towards us (and its Southeastern tip receding from us), it seems unlikely that an uneven spatial distribution of intervening material could be responsible of the observed asymmetry as it may be the case of the dusty PNe BD+30$^\\circ$3639 \\citep{KN00} and NGC\\,7027 \\citep{KN01}, because only small extinction variations were observed in NGC\\,3242 \\citep{Balick_etal93}. On the other hand, our simulations in Section 3.1 suggest that the asymmetric distribution of X-ray emission could result from small number statistics."
    },
    "1107/1107.5617_arXiv.txt": {
        "abstract": "What is the size of the \\ii{most} massive object one expects to find in a survey of a given volume? In this paper, we present a solution to this problem using Extreme-Value Statistics, taking into account primordial non-Gaussianity and its effects on the abundance and the clustering of rare objects. We calculate the probability density function (pdf) of extreme-mass clusters in a survey volume, and show how primordial non-Gaussianity shifts the peak of this pdf. We also study the sensitivity of the extreme-value pdfs to changes in the mass functions, survey volume, redshift coverage and the normalization of the matter power spectrum, $\\sigma_8$. For `local' non-Gaussianity parametrized by $\\fnl$, our correction for the extreme-value pdf due to the bias is important when $\\fnl\\gtrsim100$, and becomes more significant for wider and deeper surveys. Applying our formalism to the massive high-redshift cluster XMMUJ0044.0-2-33, we find that its existence is consistent with $\\fnl=0$, although the conclusion is sensitive to the assumed values of $\\fsky$ and $\\sigma_8$. We also discuss the convergence of the extreme-value distribution to one of the three possible asymptotic forms, and argue that the convergence is insensitive to the presence of non-Gaussianity. ",
        "introduction": "The statistics of the primordial seeds that grew into the observed large-scale structures holds a wealth of information about the physics of the primordial Universe. In the simplest models of inflation, the primordial density fluctuations obey an almost Gaussian statistics (see \\cite{bartolo} for a review). Tiny deviations from Gaussianity may be quantified, amongst other ways\\footnote{In general, a non-Gaussian pdf can have divergent moments (e.g. the Cauchy distribution). In this work we assume that the primordial density distribution has finite moments up to third order (\\ii{i.e.} finite skewness).}, by the `local' non-Gaussianity parameter, $\\fnl$, defined via the expansion of the non-linear Newtonian potential \\be \\Phi = \\phi + \\fnl(\\phi^2-\\bkta{\\phi^2})+\\ldots,\\ee where $\\phi$ is a Gaussian random field. This form of non-Gaussianity arises in simple models of single and multi-field inflation \\citep{maldacena,rigopoulos,byrnes} as well as some curvaton models \\citep{bartolo2,sasaki,malik}. Observational constraints on $\\fnl$  from the cosmic microwave background (CMB) anisotropies are currently consistent with $\\fnl=32\\pm42$ ($2\\sigma$) \\citep{komatsu}. However, if $\\fnl$ is in fact much smaller, its effects on the CMB would be difficult to extract and distinguished from non-Gaussianity arising from secondary sources such as gravitational lensing and instrumental noise \\cite{cooray,yu}.    The statistics of large-scale structures offers a complementary probe of non-Gaussianity on much smaller scales than the CMB. Indeed, it is possible that $\\fnl$ measured on Gpc scales may be quite different from that measured on Mpc scales. In wavenumber space, this translates to a possible $k$-dependence of $\\fnl$, which have been hinted at by the numerous observations of massive high-redshift clusters \\cite{jee,stott,brodwin,santos,tanaka}. These massive clusters exist, according to some, in greater abundances than expected from a Gaussian statistics. Some authors have concluded that the level of non-Gaussianity on Mpc scale required to explain the existence of certain rare clusters is $\\fnl=$ a few $\\times10^2$ \\cite{enqvist,hoyle}. In contrast, some have argued that these claims are based on misinterpretation of data, and that the occurrences of these rare objects are in fact consistent with a Gaussian statistics \\cite{hotchkiss, cayon,mortonson}.  In this work, we offer our opinion on this debate by presenting an approach to calculating the probability of observing rare objects based on extreme-value statistics. We begin by asking: what is the probability distribution of the \\ii{most} massive clusters found within a given volume at a given redshift range? Our technique relies on a basic application of the so-called void probability distribution introduced by White \\cite{white}. This approach was successfully used to study the abundances of massive clusters given a Gaussian initial condition in \\cite{davis,colombi}. In this work, we extend the groundwork laid by these authors to study the effect of $\\fnl$ on the distribution of extreme-mass objects. For other cosmological applications of extreme-value theory, see \\cite{bernstein,antal,dobos,bhavsar, mikelsons,sheth,waizmann,waizmann2}.  Previous approaches to extreme-value statistics of clusters have so far either neglected the clustering, or bias, of galaxy clusters \\cite{waizmann2, harrison}, or considered it in the Gaussian case \\cite{davis,sheth}. In this work, however, we have included the effects of the bias in the presence of non-Gaussianity. Whilst Davis \\ii{et al.} \\cite{davis} have previously reported that the effects of the bias on the extreme-value distribution are small in the Gaussian case, it remains to be shown if this also holds in the presence of non-Gaussianity, which can introduce a strong scale dependence in the bias \\cite{dalal}. We investigate this problem in this work using the formalism of Valageas \\cite{valageasA,valageasB}, who showed how the bias can be calculated in real space for a given $\\fnl$. As we shall see later, the contribution from the bias can indeed be significant if $\\fnl$ is sufficiently large. ",
        "conclusions": "In summary, we have investigated quantitatively how the statistics of extreme-mass clusters is affected by uncertainties in the mass function, non-Gaussian corrections of the mass function and bias, Eddington correction, $\\fsky$, redshift, $\\fnl$ and $\\sigma_8$. More specifically,  \\begin{enumerate} \\item We have presented a procedure to calculate the statistics of extreme-mass galaxy clusters in the presence of primordial non-Gaussianity parametrized by $\\fnl$. Our main results are the expressions for the cumulative probability distribution \\re{logp} and the probability density function \\re{pdf} for the most massive object in a survey of a given sky coverage and redshift range. These expressions enable us to deduce the most probable extreme-mass cluster in a survey of a given specification. The effects of changing the mass function and varying the value of $\\fnl$, survey volume and redshift are summarised in figures \\ref{figbig} and \\ref{varyL}.  \\item Our correction terms for the extreme-value distribution (second and third terms on the right-hand side of \\re{logp}) are significant when considering a large volume, high redshift or large non-Gaussianity (see figure \\ref{figrelative}). For non-Gaussianity with $\\fnl=\\mc{O}(1)$, the first term of \\re{logp} (Poisson approximation) suffices.  \\item Next, we applied our formalism to investigate the extreme-value properties of cluster XMMUJ0044.0-2-33 ($M\\sub{obs}\\simeq 3.12\\times10^{14}h^{-1}M_\\sun$ at $z=1.579$). We find that the probability that the cluster is the most likely extreme-mass cluster in the survey depends sensitively on the assumed sky coverage, but is consistent with $\\fnl=0$ (assuming $\\sigma_8=0.801$). More conservatively, with $\\fsky=1$, the most probable extreme-mass cluster is expected to be much larger and perhaps this will be confirmed by future X-ray cluster surveys.  \\item We show that the effect of $\\fnl$ in shifting the extreme-mass cluster to higher values is degenerate with an increase in $\\sigma_8$ (figure \\ref{figothers}). The degeneracy can be broken by combining cluster data with CMB constraints.  \\end{enumerate}   An important ingredient in our calculation is the mass function. In the presence of primordial non-Gaussianity, it remains to be seen what the correct mass function should be. Our investigation showed that the Press-Schechter, Sheth-Tormen and Tinker mass functions give similar extreme-value statistics at low redshift, but there are large differences at high redshift and large $\\fnl$. The understanding of the correct form of the mass function appropriate for these extreme-mass objects is important since the uncertainty in the distribution of extreme-mass clusters due to the mass function is comparable with that from the mass determination (typically $\\sigma_{\\ln M}\\sim0.3$). Thus, it remains for further numerical simulations along the lines of \\cite{wagner,pillepich} to establish the validity of the various mass functions and non-Gaussian correction factors in the presence of non-Gaussianity.    Another crucial ingredient is the bias which, in this work, was calculated using the real-space formalism given in \\cite{valageasA,valageasB} combined with our averaging procedure outlined in Appendix A. As pointed out in these papers, it is possible to extend the calculation to other types of non-Gaussianity (non-local or higher-order local type). It will be an interesting extension to study extreme-value statistics in the presence of different types of non-Gaussianity.        \\mmm \\sss \\bbb  \\centerline{\\bb{Acknowledgment}}  \\mmm  We are indebted to the anonymous referee, whose many insightful comments led to a major improvement of the paper. We are also grateful to Olaf Davis for helpful discussions in the initial stages, and to Christopher Gordon, Aseem Paranjape, Shaun Hotchkiss and, in particular, Jean-Claude Waizmann for their  comments on an early version of the manuscript. SC supported by Lincoln College, Oxford. \\bbb    \\appendix"
    },
    "1107/1107.5084_arXiv.txt": {
        "abstract": "{To investigate the distribution of rotation periods of asteroids from  different regions of the Solar System and distribution of diameters of near-Earth asteroids (NEAs).} {Verify if nonextensive statistics satisfactorily describes the data.} {Light curve data was taken from Planetary Database System (PDS) with Rel $\\ge 2$. Taxonomic class and region of the Solar System was also considered. Data of NEA were taken from Minor Planet Center.} {The rotation periods of asteroids follow a $q$-Gaussian with $q=2.6$ regardless of taxonomy, diameter or region of the Solar System of the object. The distribution of rotation periods is influenced by observational bias. The diameters of NEAs are described by a $q$-exponential with $q=1.3$. According to this distribution, there are expected to be $994 \\pm 30$ NEAs with diameters greater than 1~km. } {} ",
        "introduction": "Asteroids and comets are primordial bodies of the Solar System (SS). The study of the physical properties of these objects may lead to a better understanding of the processes of formation of the SS, and, by inference, of the hundreds of exo-Solar systems already known. Distribution of rotation periods and diameters of asteroids are two parameters that may give pieces of information concerning the evolution of the SS. The first attempt to describe histograms of rotation periods of asteroids was made by Harris \\& Burns (\\cite{harris-burns}). This work and also others that have followed have shown that rotation periods of big asteroids ($D>$30--40~km) follow a Maxwellian distribution. Harris \\& Pravec (\\cite{pravec-harris-2000}) have analyzed a sample with 984 objects and have confirmed that the distribution of rotation periods of asteroids with diameters $D\\ge 40$ km is Maxwellian, with 99\\% of confidence, though this hypothesis can be rejected at 95\\% of confidence. They suggest that objects within this diameter range are primordial, or originated from collisions of primordial bodies. It is known that for median sized ($10<D\\le 40$~km ) and small ($D< 10$ km) asteroids, the distribution of rotation periods is not Maxwellian. The analysis of the data suggests the existence of a spin-barrier for the asteroids with diameters between hundreds of meters and 10 km and with more than 11 rotations per day (d$^{-1}$) (period of about 2.2 h). The absence of a substantial quantity of asteroids with period less than 2.2~h may be due to the low degree of internal cohesion of these objects. The majority of the sample may contain rubble pile asteroids (Davis et al.\\ \\cite{davis-1979}, Harris \\cite{harris-1996}) that are composed by fragments of rocks kept together by self-gravitation. For objects below 0.2~km it is observed rotation periods smaller than the spin-barrier, suggesting that these objects have a high internal cohesion, implying that they may be monolithic bodies. The difficulty in the modelling of rotation periods of asteroids as a whole may be associated to the combined action of many mechanisms such as collisions (Paolicchi, Burns \\& Weidenschilling \\cite{paolicchi}), gravitational interactions with planets (Scheeres, Marzari \\& Rossi \\cite{scheeres-2004}), angular momentum exchange in binary or multiple asteroid systems (Scheeres (\\cite{scheeres-2002})), or torques induced by solar radiation, known as YORP effect (from Yarkovsky--O'Keefe--Radzievskii--Paddack) (Rubincam \\cite{rubincam}). Particularly, YORP effect is strongly dependent on the shape and size of the object and its distance to the Sun.   Near-Earth Asteroids (NEAs) is a subgroup of SS asteroids whose heliocentric orbits lead them close to the Earth's orbit. More than 7000 NEAs are known up to 2011. The study of these objects is relevant once it may bring information regarding the birth and dynamic evolution of the SS. Also a special interest in these objects is related to the possibility of collision with the Earth with obvious catastrophic consequences (Alvarez et al.\\ \\cite{alvarez-1980}). They also may be potential sources of raw material for future space projects.   The evaluation of the number of asteroids per year that may reach the Earth as a function of their diameters is essential for the determination of the potential risk of a collision. One of the first attempts to estimate this flux was done by Shoemaker et al.\\ (\\cite{shoemaker-1979}).  The impact flux may be taken from the accumulated distribution of diameters of the NEAs. This distribution is indirectly obtained by the current asteroid surveys, according to the absolute magnitude $H$. The distribution of absolute magnitude $H$ is described by Jedicke, Larsen \\& Spahr (\\cite{jedicke-2002}), \\begin{equation} \\log N = \\alpha H + \\beta, \\label{eq:distrib-h} \\end{equation} where $N$ is the number of objects, $\\alpha$ is the ``slope parameter'' and $\\beta$ is a constant. This relation asymptotically models the observed distribution of $H$. Departure from this power law is probably associated with the observational bias due to physical and dynamical properties of the asteroids (orbital elements, size, albedo), and instrumental limitations (CCD, detection software, among others). So Eq. (\\ref{eq:distrib-h}) may be used to describe a given population of asteroids if a correction of the bias is made in the raw data.  The diameters of asteroids $D$ may be given in function of their absolute magnitudes and their albedos $p_{_V}$ according to Bowell et al.\\ (\\cite{bowell}) \\begin{equation} D = 1329 \\frac{10^{-H/5}}{\\sqrt{p_{_V}}}. \\label{eq:d-h} \\end{equation} The albedo is the rate of superficial reflection and its value is essential for the estimation of the diameters of the asteroids. The values of the albedos of asteroids varies according to the superficial mineralogical composition (taxonomic complex) and object's shape. Typical values range from $0.06\\pm 0.02$ for low albedo objects of C taxonomic complex up to $0.46\\pm0.06$ for high albedo objects of V type (Warner, Harris \\& Pravec \\cite{warner-harris-pravec-2009}).  Combination of equations (\\ref{eq:distrib-h}) and (\\ref{eq:d-h}) leads to a power law behavior: \\begin{equation} N(>D)=kD^{-b}. \\label{eq:distrib-d} \\end{equation} The parameters are estimated by Stuart (\\cite{stuart-2003}), $b=1.95$ ($\\alpha=b/5$, admitting the same albedo for all the sample) $k=1090$, and $D$ is given in km. According to this expression and taking into account the uncertainties of the measures, Stuart \\& Binzel (\\cite{stuart-binzel}) have estimated that there may exist $1090\\pm 180$ objects with diameters equal to or greater than 1~km ($H=17.8$). ",
        "conclusions": ""
    },
    "1107/1107.4561_arXiv.txt": {
        "abstract": "{Over the past decade a new generation of chemical models, in addition to the gas, have included the dust in the treatment of the interstellar medium (ISM). This major accomplishment has been spurred by the growing amounts of data on the highly obscured high-z Universe and the intriguing local properties of the Solar Neighbourhood (SoNe) of the Milky Way (MW) Disk.} {We present here a new model able to simulate the formation and evolution of dust in the ISM of the MW. The model follows the evolution of 16 elemental species, with particular attention to those  that are simultaneously present in form of gas and dust, e.g. C, N, O, Mg, Si, S, Ca and Fe. In this study we focus on the SoNe and the MW Disk as a whole which are considered as laboratories to test the physical ingredients governing the dust evolution.} {The MW is described as a set of concentric rings of  which we follow the time evolution of  gas and dust. Infall of primordial gas, birth and death of stars, radial flows of matter between contiguous shells, presence of a central bar, star-dust emission by SN{\\ae} and AGB stars, dust destruction and accretion are taken into account. The model  reproduces the local depletion of the elements in the gas, and simultaneously satisfies other constraints obtained from the observations.} {The evolution of the element abundances  in the gas and dust has been well reproduced for plausible choices of the parameters. The Mg/Si ratio, in particular, drives the formation of silicates. We show that for most of the evolution of the MW, the main process for dust  enrichment is the accretion in the cold regions of the ISM. SN{\\ae} dominate in the early phases of the evolution. We have also examined the main factors controlling the temporal window in which SN{\\ae} govern the dust budget both in low and high star forming environments. The role played by AGB stars is also discussed. We find that IMFs with regular slope in the range of massive stars better reproduce the observed depletions.} {The classical chemical models nicely reproduce the abundances, depletion factors and dust properties of the SoNe and the main ingredients of the models are tested against observational data. The results obtained for the SoNe lead us to safely extend the model to the whole Galactic Disk or galaxies of different morphological types.} ",
        "introduction": "In the fascinating subject of the origin and evolution of galaxies, the interstellar dust is acquiring a primary role because of its growing importance in the observations of the high-z Universe \\citep{Omont01,Shapley01,Bertoldi03,Robson04,Wang08a,Wang08b,Gallerani10,Michalowski10a,Michalowski10b} and the theoretical spectro-photometric, dynamical, and chemical modeling of galaxies \\citep{Schurer09,Narayanan10,Jonsson10,Grassi10,Pipino11,Popescu11}. \\\\ \\indent Indeed, the evidence of highly obscured QSOs and galaxies already in place at high-z leads necessarily to a new generation of theoretical models where dust is a key ingredient that cannot be neglected, if we want to obtain precious clues on the fundamental question about when and how galaxies formed and evolved. First of all, dust absorbs the stellar radiation and re-emits it in the infrared deeply changing the shape of the observed spectral energy distributions (SEDs) of obscured galaxies \\citep{Silva98,Piovan06b,Popescu11}; second, it strongly affects the production of molecular hydrogen and the local amount of UV radiation in galaxies thus playing a strong role in the star formation process via the cooling mechanisms \\citep{Yamasawa11}. The inclusion of dust in the models leads to a growing complexity and typically to a much larger set of parameters influencing the results of the simulations to be then compared with the observations. Indeed several question must be addressed, each one easily expanding the model: who are the main stardust injectors in the interstellar medium (ISM) \\citep{Gail09,Valiante09,Gall11a,Piovan11a}? How much dust do they produce and on which timescales \\citep{Draine09,Dwek09}? What is the contribution and the role played by the molecular clouds (MCs)-grown dust that form in the cold dense regions of the ISM \\citep{Zhukovska08,Pipino11}? How much dust is destroyed by SN{\\ae} shocks \\citep{Nozawa06,Nozawa07,Bianchi07,Jones11}? What is the typical minimal set of dust grains whose evolution should be followed and what could be a minimal set of dust grains to be used for a satisfactory description of the chemical or spectrophotometric properties of the galaxy? \\\\ \\indent To answer all these questions the theoretical models must include (i) a set of grains with suitable composition and properties and/or an ISM made of gas and dust in which the abundances of the elements are followed, (ii) a recipe for their formation/accretion and destruction in the ISM and (iii) a prescription for the yields of dust by the stellar sources. The duty cycle of the dust can be schematically summarized as follows \\citep{Jones04}. Stars, mainly AGBs and SN{\\ae}, inject material in the ISM, mainly in form of gas, but with a variable amount that condenses into the so-called star-dust. Once injected into the ISM, star-dust grains are subjected to destruction processes that restitute the material to the gaseous phase. However in cold and dense regions dust can accrete on the so-called seeds: the competition between the accretion and destruction processes, mainly via shocks, determines the total budget of dust in the ISM and the observed depletion of the elements that are involved in the formation of dust grains \\citep{Dwek98}. Dust accretion mainly occurs in the very cold molecular clouds (MCs), where it induces strong cooling thus leading to the formation of new stars. The stellar winds from AGB stars SN{a} explosions more and more enrich the ISM with new metals and star-dust grains that are able to survive to the local shocks caused by  SNa explosions. \\\\ \\indent The Milky Way (MW) is the ideal laboratory to our disposal to study the dust cycle  \\citep{Zhukovska08} and its impact on the wider subject of galaxy formation. For obvious reasons, the MW provides plenty of observational data to which we can compare theoretical predictions, thus setting useful constraints on theoretical simulations and highlighting the role of the most important physical quantities leading the whole problem. Once this important step is accomplished, our modelling of the role played by dust can be extended to other galaxies such as local disk and spheroidal galaxies, high-z galaxies and QSOs. \\\\ \\indent Starting from these considerations, in this paper we simulate the formation, evolution and composition of dust in the MW, both locally in the Solar Neighbourhood (SoNe) and radially along the Galactic Disk. We build up a detailed chemical model (the theoretical simulation is still the  main tool to investigate the formation and evolution of dust in galaxies) starting from the pioneering study  by \\citet{Dwek98} and taking into account the more recent ones by \\citet{Zhukovska08}, \\citet{Calura08}, \\citet{Valiante09}, \\citet{Gall11a}, \\citet{Gall11b}, \\citet{Mattsson11}, \\citet{Valiante11}, \\citet{Kemper11} and \\citet{Dwek11}. The theoretical model we are building up stems from the basic one with infall by \\citet{Chiosi80}, however updated to the more recent version with radial flows of matter and presence of a central bar  developed by \\citet{Portinari00}. The stellar yields of chemical elements in form of gas  are those calculated by \\citet{Portinari98}. The model follows the evolution of the abundances of  a number of elements composing the ISM gas, includes the formation/destruction and evolution of dust and, finally, follows in detail also the abundances of those elements that are embedded in the dust grains. To this aims, the model makes use of the best prescriptions available in literature concerning dust accretion, destruction and condensation in the AGB  winds and SN{\\ae} explosions.  These prescriptions have been already presented by \\citet{Piovan11a} to whom the reader should refer and will also be discussed in some detail here.  \\\\ \\indent The main test for any  model of dust formation is given by the data on element depletion provided by SoNe of the MW to which we will compare our results. In a forthcoming paper \\citep{Piovan11c} we will investigate the radial dependence of chemical abundances and dust depletion across the Disk of the MW. \\\\ \\indent The plan of the paper is as follows. In Sect. \\ref{Che_Evo_Mod} we introduce the formalism and basic equations governing the temporal evolution of the gas, stars, and dust in an open model with radial flows of gas and dust for the galactic Disk of the MW. In Sect. \\ref{SFRandIMF} we summarize the current prescriptions for the star formation rate and initial mass function. In Sect. \\ref{graingrowth} we introduce and describe in some detail the various processes responsible for the formation and growth of dust grains in the ISM, whereas in  Sect. \\ref{grainevolution} we present the accretion rates into dust grains for a number of important elements we have considered. The yields of dust from AGB stars and SNa explosions are adopted according to \\citet{Piovan11a} to whom the reader should refer for all the details. In Sect. \\ref{SNEdestruction} the problem of dust destruction by SN{a} shocks is faced. Then, in Sect. \\ref{AbundancesSS} we present the observational data for the elemental abundances in the Solar Neighborhood and define the reference set of abundances we have adopted. In Sect. \\ref{Dustymodels} after summarizing the main ones between the many available parameters, we discuss and compare the effect of them on the formation and evolution of dust, at varying them between the many possible choices. In particular we examine the influence on dust of a different CO fraction in the ISM (Sect. \\ref{COFraction}), the effect of the IMF (Sect. \\ref{IMFeffect}) and of the SF law (Sect. \\ref{SFeffect}), the choice between different models for the accretion of dust in cold regions (Sect. \\ref{Accreffect}) and, finally we discuss some interesting parameters however not discussed in our model. In Sect. \\ref{Models_Observ} we present the results for our models of the Solar Vicinity in presence of dust and as a function of three important ingredients, namely, the initial mass function, the efficiency of star formation, the accretion time scale of primordial gas onto the system mimicking the evolution of the MW Disk. The effect of radial flows and central bar are always included according to the prescription developed in previous studies \\citep{Portinari00} and also adopted in \\citet{Piovan11c}. The simulations are compared with the depletion of the elements in the SoNe under the constraints that we want that the local chemical properties are satisfied, such as the time evolution of the elemental abundances, the metallicity and iron enrichment. In Sect. \\ref{Discus_Concl} we discuss the results we have obtained and draw some general conclusions. ",
        "conclusions": "In the first paper of this series of three \\citep{Piovan11a}, we presented and discussed the prescriptions currently in use to describe the type and amounts of dust injected into the ISM by AGB stars and SNa explosions.  The  condensation coefficients of the dust have tested and analyzed referring to a suitable  chemical model for the MW Disk and the Solar Neighbourhood in particular. This reference model and its physical ingredients are described in great detail in this paper.  In particular, we focused on the mechanism of dust accretion in the ISM,  dust destruction by various processes, the main parameters of the chemical model, namely IMF, star formation efficiency and infall timescale. The main conclusions can be summarized as follows:  \\begin{itemize} \\item The CO molecules influence the formation of C-based dust, thus introducing into the estimates an unavoidable uncertainty. The higher the amount of carbon embedded into CO, the slower is the accretion process and the smaller the amount of carbonaceous grains that are  formed.  \\item During most of the MW evolutionary history,  the main process enriching the ISM in dust is the accretion in the cold regions. Only in the very early stages,  SN{\\ae} dominate, and the duration of this phase tends to shorten in regions of  high star formation and fast enrichment in metals (innermost regions of the MW). The opposite for the regions of low star formation (the outskirts of the MW), where the accretion in the ISM becomes significant much later so that for many Gyr SN{\\ae} govern the total dust budget. In this case, AGB stars play an important  role, because there is enough time for them to significantly contribute to the  dust budget, without being overwhelmed by the SN{\\ae} (earlier phases) or the ISM accretion (later phases). The time interval during which SN{\\ae} are the main dust producer can slightly change at varying the upper mass limit of  stars undergoing  the AGB phase.  Our conclusions for the MW Disk could be extended to  galaxies characterized by  continuous star formation on the notion that the outer regions of the MW Disk might correspond to low star forming galaxies and the  inner ones to high star forming objects. Very high values of SFR \\citep[See for example] [for detailed simulations of starburst galaxies and QSOs]{Gall11a,Gall11b}, which are not reached in the MW, not even in the early phases of the evolution in the inner regions, somehow elude this simple scheme. In any case, star-dust dominates the mild SF environments for a long  period of time. This time scale tends to decrease in environments with IMFs skewed toward the mass interval in which  the star-dust/metals factories (AGB stars and SN{\\ae}) are important, thus helping the onset of the accretion phase in the ISM.   \\item In the high SFR/high metallicity regions AGB stars mainly produce silicates, whereas in the low SFR/low metallicity ones, carbon stars can contribute significantly to the C-based dust.  \\item The IMF plays a fundamental role because it controls the relative amounts of low, intermediate and massive stars (see the entries of Table \\ref{IMFfractions}). We found that IMFs skewed toward massive stars are not suitable to reproduce the properties of dust in the  MW, the SoNe in particular, \\citep[see also][]{Dwek98,Zhukovska08,Calura08}, whereas it seems that for star-burst galaxies \\citep{Gall11a}, ellipticals \\citep{Pipino11} and QSOs \\citep{Gall11b,Valiante11} IMFs more biased toward higher mass stars are required, because otherwise it would be more difficult  to reproduce the amount of dust observed already on site in very high-z objects. The alternative is to introduce very high condensation efficiencies in SN{\\ae}. This possibility is somehow supported by the recent FIR/sub-mm observations on the SN 1987A by \\citet{Matsuura11}. In such a case normal  IMFs in high-z galaxies and QSOs cannot be excluded. The use of different IMFs for different galaxies in order to reproduce dust properties is a subject of vivid debate, and tightly related to the wider question about the universality of the IMF. The debate seems to favour an IMF sensitive to the initial conditions of star formation \\citep[see for instance the recent] [\\,to mention a few]{Elmegreen09,Bastian10,Myers11,Gunawardhana11,Kroupa11}. This possibility was suggested long ago by \\citet{Chiosi1998} to solve the apparent contradiction between the spectral and chemical properties of early type galaxies. Similar conclusions are reached by \\citet{Valiante11}, trying to match the observed properties of the QSOs SDSS J1148+5251, where a top-heavy IMF allow a more coherent match between different observations. Finally, when the SFR is high (inner regions), the effect of the IMF on the dust budget tends to disappear at the current time, because the accretion in the ISM becomes dominant early on  in the evolution, while if the SFR is low we can see the effect of different IMFs spreading until the current time.  \\item Lower efficiencies $\\nu$ of the SF, in the range adopted for the SoNe, correspond to a slower onset of the accretion in the ISM and a lower final dust budget.  \\item We tested different descriptions of the process of dust accretion in the ISM. A simple approach adopts an average time-scale of accretion to estimate the \\textit{total} dust budget at least in the early stages and for normal star forming environments. More complicated descriptions are require to  follow the evolution of the abundance of single elements in dust or low star forming regions.  \\item The range of depletions observed for most of the elements, that is C, N, O, Mg, Si and S, is nicely  reproduced by  reasonable combinations of the parameters. IMFs with the classical slope in the high mass range,  more easily fit  the data for depletion. The only exceptions are: (i) iron, for which probably a more complex mechanism of accretion in the ISM than the simple one on cold regions is  required \\citep{Zhukovska08}; (ii)  calcium whose extreme depletion can not be reproduced by our model.  \\item The ratio between the abundances of Mg and Si is crucial in the formation of silicates and it is important to tune it according to the ratio expected from models.  We had to  correct  the slightly under-abundance of the Mg in our gaseous yields, to obtain  a better agreement with the observed depletion for the refractory elements involved into the silicates formation. \\end{itemize}  To conclude, the classical chemical models nicely reproduce the observed depletion and properties of the SoNe. The theoretical ingredients, like condensation coefficients and accretion models, behave in a satisfactory way. What can be improved? Clearly the major weak point is the one-phase description of the ISM. At least a two-phases ISM is required. Finally, work is in progress to introduce a multi-phase description of the ISM with dust in N-Body-TSPH simulations of galaxies.\\\\  \\textit{Acknowledgements}. L. Piovan acknowledges A. Weiss and the Max Planck Institut F\\\"ur AstroPhysik (Garching - Germany) for the very warm and friendly hospitality and for providing unlimited computational support during the visits as EARA fellow when a significant part of this study has been carried out. The authors are also deeply grateful to  S. Zhukovska and H. P. Gail for many explanations and clarifications about their model of dust accretion, T. Nozawa and H. Umeda for many fruitful discussions about SNa dust yields. This work has been financed by the University of Padua with the dedicated fellowship \"Numerical Simulations of galaxies (dynamical, chemical and spectrophotometric models), strategies of parallelization in dynamical lagrangian approach, communication cell-to-cell into hierarchical tree codes, algorithms and optimization techniques\" as  part of the AACSE Strategic Research Project.  \\begin{scriptsize}"
    },
    "1107/1107.1728_arXiv.txt": {
        "abstract": "\\noindent We report the discovery of a serious bias in galaxy photometry reported in the 2MASS Extended Source Catalog (Jarrett \\etal 2000).  Due to an undetermined flaw in the 2MASS surface photometry routines, isophotal and total magnitudes calculated by their methods underestimate the luminosity of galaxies from 10\\% to 40\\%.  This is found to be due to incorrectly determined scalelengths and isophotal radii, which are used to define the aperture sizes for Kron and total fluxes.  While 2MASS metric aperture luminosities are correct (and, thus, colors based on those apertures), comparison to other filters (e.g. optical) based on total magnitudes will produce erroneous results.  We use our own galaxy photometry package (ARCHANGEL) to determine correct total magnitudes and colors using the same 2MASS images, but with a more refined surface brightness reduction scheme. Our resulting colors, and color-magnitude relation, are more in line with model expectations and previous pointed observations. ",
        "introduction": "Surface photometry is an important tool in the understanding of galaxy mass and structure.  Galaxy luminosity measures the primary baryonic component, (i.e. stellar mass) and structural information traces the gravitational potential.  Galaxy formation scenarios make specific predictions on the light distribution of galaxies, so accurate reduction of a galaxy image into a total magnitude and scalelength are important parameters to understanding the fundamental plane and the star formation history of galaxies.  Obtaining the structural characteristics and the total luminosity of a galaxy requires knowledge of its surface brightness profile to a significant depth.  In order to extrapolate a total luminosity, isophotal analysis is required to determine how far one needs to integrate a galaxy's light plus to provide sufficient information to extrapolate the light profile.  During a surface photometry project to explore the structure of galaxies by morphological type (Schombert \\& Smith 2011), we discovered a significant discrepancy between the structural and luminosity parameters that we determined using raw 2MASS images versus those reported by the 2MASS project in their Extended Source Catalog (Jarrett \\etal 2000).  We use this letter to outline the problem, and the solution, for other researchers. ",
        "conclusions": "First, we note that this discrepancy has no impact on projects which use 2MASS aperture colors.  For galaxy colors are calculated using 2MASS total magnitudes still use the same sized apertures for $J$., $H$ and $K$, and the colors will remain consistent (although for a smaller portion of the total galaxy light).  However, comparison between other total magnitudes (e.g. RC3 magnitudes) and 2MASS Kron or total magnitudes will be biased towards the blue.  An example of this is shown in Figure 7, a histogram of $V-J$ colors for the 421 ellipticals in our sample.  As can be seen, the 2MASS colors are 0.25 mags bluer than colors calculated from our total magnitudes since the RC3 $V$ magnitude is determined from an asymptotic fit and, therefore, contains more flux that 2MASS's total magnitude.   There appears to be no standard correction from 2MASS colors to the correct colors, this would require information on how deviant the 2MASS surface brightness profiles (from which the aperture sizes are extracted) are from reality.  There appears to be no simple formula.  A priority science goal for 2MASS was large baseline color comparison.  An example of relevance of large wavelength comparisons is the color-magnitude relation (CMR).  The CMR is a long known correlation between galaxy color and luminosity.  The best explanation is that galaxies with higher mass have higher metallicities.  Global metallicity reflects in the mean temperature of the RGB such that low metallicities produce bluer colors. The CMR for this data sample is shown in Figure 8.  Again, we see that the 2MASS colors predict the {\\it opposite} expectation from earlier optical work in that they find roughly bluer colors with higher luminosity.  Using our total magnitudes (combined with RC3 colors) restores the correct CMR, redder colors with higher galaxy luminosity.  We summarize our findings as the following:  \\begin{description}  \\item{(1)} 2MASS aperture magnitudes are reproducible with high accuracy when the same metric apertures are used regardless of shape.  \\item{(2)} Comparison between our photometry packages total or Kron magnitudes and 2MASS's values (from the Extended Source Catalog, Jarrett \\etal 2000) find serious bias in the 2MASS luminosities of between 10 to 40\\%.  \\item{(3)} Point by point comparison of the surface brightness profiles provided by Jarrett \\etal (2003) indicate that problem lies in the intensity values deduced by each photometry pipeline.  2MASS values are increasingly fainter with increasing radius, this results in surface brightness profiles that are smaller in isophotal radius compared to our profiles.  \\item{(4)} Comparison to raw intensity values (e.g., the 100 arcsecs annulus in NGC 1407) demonstrates that, for unknown reasons, the 2MASS project extracts lower intensity values than given by the images.  This is either a flaw in the processing pipeline or an error in the data storage system.  \\item{(5)} Underestimated intensity values produces smaller isophotal radii at constant surface brightness and smaller scalelength fits.  Both result in smaller apertures and underestimated luminosities.  \\item{(6)} While logically, either our photometry or 2MASS's is in error, we consider three pieces of evidence to indicate our values are correct. First, the direct comparison of intensity values from the raw images. Second, the final $V-J$ colors where our colors are in agreement with mean elliptical colors in the literature, the 2MASS project's values are 0.25 too blue.  Third, there is a negative slope in the CMR from 2MASS data, ours data correct matches the color slope expected from SED models and optical CMR's.  \\end{description}  It is highly confusing on why this discrepancy has not been noticed to date.  Either we have become too trusting of electronically available datasets. Or we simply, due to pressures to publish, are not as thorough in our data analysis, especially with respect to external checks.  We becoming increasingly vulnerable to systematic errors with increasing automation to our datasets, leading to potentially highly embarrassing errors in our science results.  \\noindent Acknowledgements: Most of the comparative values were extracted from NED (NASA's Extragalactic Database) using new network tools.  The quick access to difference galaxy catalogs on one site made this project doable in reasonable timescales.  The model for future science is not faster cycles, but faster and clearer access.  The software for this project was supported by NASA's AISR and ADP programs."
    },
    "1107/1107.0775_arXiv.txt": {
        "abstract": "Radio properties of supernova outbursts remain poorly understood despite longstanding campaigns following events discovered at other wavelengths. After $\\sim 30$ years of observations, only $\\sim 50$ supernovae have been detected at radio wavelengths, none of which are Type Ia. Even the most radio-loud events are $\\sim 10^4$ fainter in the radio than in the optical; to date, such intrinsically dim objects have only been visible in the very local universe. The detection and study of radio supernovae (RSNe) will be fundamentally altered and dramatically improved as the next generation of radio telescopes comes online, including EVLA, ASKAP, and MeerKAT, and culminating in the Square Kilometer Array (SKA); the latter should be $\\ga 50$ times more sensitive than present facilities. SKA can repeatedly scan large ($\\ga 1 \\, \\rm deg^2$) areas of the sky, and thus will discover RSNe and other transient sources in a new, automatic, untargeted, and unbiased way. We estimate SKA will be able to detect core-collapse RSNe out to redshift $z \\sim 5$, with an all-redshift rate $\\sim 620 \\ {\\rm events \\ yr^{-1} \\ deg^{-2}}$, assuming a survey sensitivity of 50 nJy and radio lightcurves like those of SN 1993J. Hence SKA should provide a complete core-collapse RSN sample that is sufficient for statistical studies of radio properties of core-collapse supernovae. EVLA should find $\\sim 160 \\ {\\rm events \\ yr^{-1} \\ deg^{-2}}$ out to redshift $z \\sim 3$, and other SKA precursors should have similar detection rates. We also provided recommendations of the survey strategy to maximize the RSN detections of SKA. This new radio core-collapse supernovae sample will complement the detections from the optical searches, such as the LSST, and together provide crucial information on massive star evolution, supernova physics, and the circumstellar medium, out to high redshift. Additionally, SKA may yield the first radio Type Ia detection via follow-up of nearby events discovered  at other wavelengths. ",
        "introduction": "\\label{sect:intro}  Supernovae are among the most energetic phenomena in the universe, and are central to cosmology and astrophysics. For example, core-collapse supernovae are explosions arise from the death of massive stars and hence are closely related to the cosmic star-formation rate and to massive-star evolution, and are responsible for the energy and baryonic feedback of the environment \\citep{madau}. Type Ia supernovae show a uniform properties in their lightcurves and play a crucial role as cosmic ``standardizable candles'' \\citep{Phillips93,Riess96}.  Our knowledge of the {\\em optical} properties of supernovae, is increasing rapidly with the advent of prototype ``synoptic''--i.e., repeated scan--sky surveys, such as SDSS-II \\citep{Frieman08,Sako08} and SNLS \\citep{Bazin09,PD10}. These campaigns are precursors to the coming ``Great Survey'' era in which synoptic surveys will be conducted routinely over very large regions of sky, e.g.,~LSST \\citep{lsst,lsstsb} and Pan-STARRS \\citep{pan-starrs-sne}. The number of detected supernovae will increase by several orders of magnitude in this era \\citep{lsstsb,lf}.  In contrast to this wealth of optical information, properties of supernovae in the radio remain poorly understood, fundamentally due to observational limitations. Radio supernovae (RSNe) have primarily been discovered by follow-up observations of optical outbursts, and only very rarely by accident. To date, only $\\sim 50$ core-collapse outbursts have radio detections, and no Type Ia explosion has ever been detected in the radio \\citep{Weiler04, Panagia06}. The core-collapse subtype Ibc has been a focus of recent study in the radio, because some Type Ibc events are associated with long Gamma-Ray Bursts (GRBs) \\citep{Galama98, Kulkarni98, Soderberg07, Berger03}.  Current radio interferometers are scheduled primarily around targeted observations proposed by individual principal investigators. This stands in contrast to future radio interferometers planned for the coming ``Great Survey'' era. These include the Square Kilometer Array (SKA\\footnote{http://www.ska-telescope.org}) and its precursor prototype arrays (for example, ASKAP\\footnote{http://www.atnf.csiro.au/projects/askap} and MeerKAT\\footnote{http://www.ska.ac.za/meerkat}). These telescopes will operate primarily as wide-field survey instruments focusing on several key science projects \\citep{Carilli04}. As synoptic telescopes, they will be far better suited to study all classes of transient and time-variable radio sources, including RSNe. \\citet{Gal-Yam06} already pointed out the power of synoptic radio surveys for detecting radio transients of various types, including supernovae and GRBs, in an unbiased way. Here we quantify the prospects for RSNe.  In this paper we explore this fundamentally new mode of {\\em untargeted} RSN discovery and study. We adopt a forward-looking perspective, and consider the new science enabled by RSNe observations in an era in which the full SKA is operational. Our focus is mainly on core-collapse supernovae, the type for which some radio detections already exist. However, we will also discuss the possibility of Type Ia radio discovery based on current detection limits. We will first summarize current knowledge of radio core-collapse supernovae (\\S \\ref{sect:properties}), and the expected sensitivity of SKA (\\S \\ref{sect:sensitivity}). Using this information, we forecast the radio core-collapse supernovae harvest of SKA (\\S \\ref{sect:formalism}), and consider optimal survey strategies (\\S \\ref{sect:recommend}). We conclude by anticipating the RSN science payoff in this new era (\\S \\ref{sect:science}). We adopt a standard flat $\\Lambda$CDM model with $\\Omega_{\\rm m} = 0.274$, $\\Omega_{\\Lambda}=0.726$, and $H_0 = 70.5 \\, \\rm km \\, s^{-1} \\, Mpc^{-1}$ \\citep{wmap5} throughout. ",
        "conclusions": "\\label{sect:science}  SKA's capability for unbiased synoptic searches over large fields of view will revolutionize the discovery of radio transients in general and core-collapse RSNe in particular \\citep{Gal-Yam06}. The unprecedented sensitivity of SKA could allow detection of core-collapse RSNe out to a redshift $z \\sim 5$. These detections will be {\\em unbiased} and {\\em automatic} in that they can occur anywhere in the large SKA field of view without need for targeting based on prior detection at other wavelengths. With SKA, the core-collapse RSN inventory should increase from the current number of several dozen to $ \\sim 620 \\ \\rm yr^{-1} \\ deg^{-2}$. EVLA should detect $ \\sim 160 \\ \\rm yr^{-1} \\ deg^{-2}$, and other SKA precursors should reap similarly large RSN harvests. In contrast, intrinsically dim RSNe such as Type Ia events and 1987A-like core-collapse explosions are unlikely to be found blindly. However, the SKA (and precursor) sensitivities will offer the possibility of detecting these events via {\\em targeted} followup of discoveries by optical synoptic surveys such as LSST.  The science payoff of large-scale RSNe searches touches many areas of astrophysics and cosmology. We conclude with examples of possible science applications with the new era of RSN survey. However, the true potential of untargeted radio search is very likely beyond what we mention.  Non-prompt RSN emission requires the presence of circumstellar matter, so such surveys will probe this material and the pre-supernova winds producing it. For core-collapse supernovae, pre-supernova winds should depend on the metallicities of the progenitor stars \\citep{Leitherer92,Kudritzki00,Vink01,Mokiem07}, and should be weaker in metal-poor environments with lower opacities in progenitor atmospheres. This effect should lead to correlations between RSN luminosity and host metallicity, as well as an evolution of the RSN luminosity function towards lower values at higher redshifts. For Type Ia supernovae, the mass-loss rate from the progenitors depends on the nature of the binary system, i.e., single or double degenerate \\citep{Nomoto84,Iben84,Webbink84}. Radio detection of Type Ia supernovae will probe the mass and density profile of the surrounding environment and hence be valuable for studying Type Ia physics \\citep{Eck95,Panagia06,Chomiuk11}.  Large-scale synoptic RSNe surveys will complement their optical counterparts. While optical surveys such as LSST will provide very large supernova statistics at $z \\lesssim 1$, radio surveys will be crucial for detections at higher redshifts. The nature and evolution of dust obscuration presents a major challenge for optical supernova surveys and supernova cosmology. Current studies suggest dust obscuration increases rapidly with redshift, but uncertainties are large. \\citet{Mannucci07} estimate that optical surveys may miss $\\sim 60\\%$ of core-collapse supernovae and $\\sim 35\\%$ of Type Ia supernovae at redshift $z \\sim 2$. RSN observations, in contrast, are essentially unaffected by dust. Thus, high-redshift supernovae could be detected at radio wavelengths but largely missed in counterpart optical searches. Comparing supernova detections in both optical and radio will provide a new and independent way to measure dust dependence on redshift. In particular, SKA will be a powerful tool to directly detect supernovae in dust-obscured regions at large redshift, and therefore offer what may be the only means to study the total supernovae rate, star-formation, and dust behavior in these areas.  Additionally, radio surveys will reveal rare and exotic events. For example, some Type Ibc supernovae are linked to long GRBs \\citep[Galama et al.~1998; and see reviews in][]{Woosley06,Gehrels09}, probably via highly relativistic jets powered by central engines and will manifest themselves in extremely luminous radio emission \\citep{Woosley93, Iwamoto98, Li99, Woosley99, Heger03}. Thus one might expect radio surveys to preferentially detect more Type Ibc supernovae than other supernova types. An unbiased sample of Type Ibc RSNe will provide new information about the circumstellar environments of these explosions and thus probe the mass-loss effects believed to be crucial to the Ibc pre-explosion evolutionary path \\citep{Price02,Soderberg04,Soderberg06,Crockett07,Wellons11}; in addition, a large sample of Ibc RSNe will allow systematic study of the differences, if any, between those which do an do not host GRBs \\citep{Berger03,Soderberg-Nakar06,Soderberg07}.  Furthermore, radio surveys give unique new insight into a possible class of massive star deaths via direct collapse into black holes, with powerful neutrino bursts but no electromagnetic emission \\citep{MacFadyen99,Fryer99,MacFadyen01,Heger03}. These ``invisible collapses'' can be probed by comparing supernovae detected electromagnetically and the diffuse background of cosmic supernova neutrinos \\citep[and references therein]{lfb}. By revealing dust-enshrouded SNe, radio surveys will make this comparison robust by removing the degeneracy between truly invisible events and those which are simply optically obscured. Indeed, direct collapse events without explosions but with relativistic jets are candidates for GRB progenitors. A comparison among RSNe, optical supernovae, GRBs, and neutrino observations will provide important clues to the physics of visible and invisible collapses, and their relation with GRBs.  We thus believe that a synoptic survey in radio wavelengths will be crucial in many fields of astrophysics, for it will bring the first complete and unbiased RSN sample and systematically explore exotic radio transients. SKA will be capable of performing such an untargeted survey with its unprecedented sensitivity. Our knowledge of supernovae will thus be firmly extended into the radio and to high redshifts."
    },
    "1107/1107.0858_arXiv.txt": {
        "abstract": "This paper continues our development of non-parametric tests for analysing the completeness in apparent magnitude of magnitude-redshift surveys. The purpose of this third and final paper in our completeness series is two-fold: firstly we explore how certain forms of incompleteness for a given flux-limited galaxy redshift survey would manifest themselves in the ROBUST {\\tctv} completeness estimators introduced in our earlier papers; secondly we provide a comprehensive error propagation for these estimators.  This work was initiated by \\cite{Rauzy:2001} and then extended and developed by \\citet*{Johnston:2007} ({\\it Completeness I}) and \\citet*{Teodoro2010MNRAS.405.1187T} ({\\it Completeness II}). Here we seek to consolidate the ideas laid out in these previous papers.    In particular our goal is to provide for the observational community statistical tools that will be more easily applicable to real survey data.  By using both real surveys and Monte Carlo mock survey data, we have found distinct, characteristic behaviour of the {\\tctv} estimators which identify incompleteness in the form of e.g. missing objects within a particular magnitude range.  Conversely we have identified signatures of `over' completeness, in cases where a survey contains a small region in apparent magnitude that may have too many objects relative to the rest of the data set.  Identifying regions of incompleteness (in apparent magnitude) in this way provides a powerful means to e.g. improve  weighting schemes for estimating luminosity functions, or for more accurately determining the selection function required to employ measures of galaxy clustering as a cosmological probe.  We also demonstrate how incompleteness resulting from  luminosity evolution can be identified and provide a framework for using our estimators as a robust tool for constraining models of luminosity evolution.  Finally we explore the error propagation for \\tctv.  This builds on {\\it Completeness II} by allowing the definition of these estimators, and their errors, via an adaptive procedure that accounts for the effects of sampling error on the observed distribution of apparent magnitude and redshift in a survey. ",
        "introduction": "The notion that we have entered an era of {\\em precision cosmology\\/} has been consistently cited in cosmology publications for more than a decade (see \\citep{Turner:1998}). While many would support this statement when applied to CMBR measurements e.g. made by the Wilkinson Microwave Anisotropy Probe  \\citep[WMAP -- see e.g.][]{Spergel:2003, Larson:2010}, in other areas, such as the study of galaxy redshift surveys, the era of precision cosmology would appear to be approaching but will require improvement in both the quality and size of our datasets and (crucially) our statistical toolbox before we can claim that it has truly arrived.  Estimating the luminosity function (LF)  remains a powerful and popular probe of galaxy evolution \\citep[e.g.][]{Norberg:2002b,Blanton:2003ApJ...592..819B,Liske:2003,Richards:2006,Willmer:2006ApJ...647..853W,Faber:2007,2007MNRAS.381.1548K, Bouwens:2008,Siana:2008,Crawford:2009,Brusa:2009ApJ...693....8B,Zucca:2009,Croom:2009,Haberzettl:2009,Montero:2009MNRAS.399.1106M,Willott:2010,2010MNRAS.406.1944W,Aird:2010MNRAS.401.2531A,Rodighiero:2010AA...515A...8R}. However, its accurate determination remains one of the most fundamental statistical challenges in modern observational cosmology.  The methodology employed to estimate the LF varies greatly throughout the literature. Some of these approaches have been non-parametric, i.e. adopting no specific parametric model for the LF, and have ranged historically from the classical number count test \\citep[e.g][]{Hubble:1936,Christensen:1975}, the \\cite{Schmidt:1968} $1/V_{\\max}$ estimator, the $\\phi/\\Phi$ method \\citep[e.g.][]{Turner:1979}, and the \\cite{lynden:1971} $C^-$ method.  Alternatively, there have also been several parametric methods developed, generally based on the Maximum Likelihood Estimator (MLE) method of \\citep{sandage:1979} (STY), where a parametric form  of the LF  \\citep[most commonly that of][]{Schechter:1976} is assumed. Thirdly,  there exists a `hybrid' method: the non-parametric counterpart of the MLE developed by \\cite{Efstathiou:1988}, and often referred to as the Stepwise Maximum Likelihood method (SWML).   Of course we have not listed the many variations of these three broad approaches that have arisen as a result of a now thriving industry that has produced (and continues to produce) the myriad of complex and tailor-made galaxy redshift surveys.  For more detailed articles that trace the origins and development of this area of statistics the reader is referred to  e.g. \\cite{2011AARv..19...41J,Binggeli:1988,Willmer:1997} and \\cite{Takeuchi:2000}  Recently the field of LF estimation has witnessed some significant fresh developments, where new approaches have emerged with the goal of placing the methodology on a more rigorous statistical footing. Examples of these include: the semi-parametric  approach  by \\cite{Schafer:2007}; Bayesian methods of   \\cite{Andreon:2006MNRAS.369..969A} and \\cite{Kelly:2008},  the use of the  {\\it copula}\\footnote{The copula  is a function used to  {\\it join} multivariate distribution functions to their one-dimensional marginal distribution function and is particularly useful for variables with co-dependence.}   by  \\cite{Takeuchi2010MNRAS.406.1830T} to   construct the bivariate LF; and a non-parametric inversion technique by \\cite{Borgne:2009AA...504..727L} applied  to galaxy counts.  Two key  fundamental assumptions that are common to almost all LF estimators and which are also crucial to the work detailed in this paper can be summarised as: \\begin{enumerate} \\item{{\\it Separability} between the luminosity function, $\\phi (M)$,  and the density function, $\\rho(z)$, probability densities; i.e. one assumes that the underlying joint distribution of luminosity (equivalently absolute magnitude) and redshift may be written as the product of their marginal distributions.} \\item{{\\it Completeness} in apparent magnitude up to a specified faint apparent magnitude limit or within specified bright and faint magnitude limits.} \\end{enumerate} For clarity, note that we define completeness in this context as the probability that a galaxy of apparent magnitude, $m$, is observable.  In \\cite{Johnston:2007}  (hereafter  {\\CI}) and \\cite{Teodoro2010MNRAS.405.1187T}  (hereafter {\\CII}) we discussed in some depth the relative merits of traditional apparent magnitude completeness tests -- e.g. the classical number count test \\citep{Hubble:1926} and the still widely used \\cite{Schmidt:1968} {\\vmax} test.  In particular we highlighted that both tests are susceptible to bias when applied to a survey which is spatially inhomogeneous.  More specifically, we noted that in practice it can be difficult to decide whether deviations from the expected value of the respective test statistics are indeed the result of survey incompleteness in apparent magnitude, or are an artefact of galaxy clustering and/or evolution of the galaxy luminosity function.   However, at least in the case of {\\vmax} (and its $1/V_{\\max}$ counterpart), this issue has  not diminished the widespread application and extension of these tests -- possibly due to their simple implementation \\citep[see .e.g.][]{Huchra:1973,Felten:1976,avni:1980,Hudson:1991,Eales:1993,Qin:1997,Qin:1999,Page:2000,Sheth:2007}.  Nevertheless, we believe that this should not deter us from developing better statistical tests that are less prone to such biases when the assumptions underpinning them are not fully satisfied.  \\cite{Efron:1992} (hereafter EP92) revisited the properties of the  seminal \\cite{lynden:1971}  $C^-$ method for constructing galaxy LFs.  Their paper introduced a powerful new approach to analysing magnitude-redshift surveys in which they proposed a non-parametric permutation test of the independence of the spatial and luminosity distributions of galaxies in a magnitude-limited sample.  As with the $C^-$ method, the EP92 test  required no assumptions concerning the parametric form of both the spatial distribution and the galaxy luminosity function.  However, this estimator  assumes the data under test {\\it is\\/} complete in apparent magnitude. They applied their permutation test to a quasar sample, with an assumed apparent magnitude limit, in order to robustly estimate the parameters characterising the luminosity distance-redshift relation of the quasars \\citep[see also][]{Efron:1998}.  These permutation test tools have since been adopted to explore e.g. evolutionary  models for the dependence of Gamma Ray Burster (GRB) luminosity with redshift \\citep{Lloyd:2002}, correlations between the luminosity of galactic nuclei and that of their host galaxy \\citep{Hao:2005}, and, more recently, the radio and optical luminosity functions in \\cite{Singal:2011arXiv1101.2930S}.  \\cite{Rauzy:2001} (hereafter R01) extended the ideas of EP92  by adapting them to develop a simple but powerful tool for assessing the apparent magnitude completeness of magnitude-redshift surveys. As was the case with EP92 -- and unlike the Hubble number counts or {\\vmax} tests -- the Rauzy test statistic {\\tc} requires {\\it no\\/} assumption that the spatial distribution of galaxies is homogenous.  Moreover, it also requires no assumption of a specific parametric form for the galaxy luminosity function. However, the Rauzy test was formulated only for the case of an assumed sharp, faint apparent magnitude limit. The Rauzy test has since been applied to a wide variety of data exploring e.g.  the completeness limits of the HI mass function in the HIPASS survey  \\citep{Zwaan:2004MNRAS.350.1210Z},  the HI flux completeness of ALFALFA survey data \\citep{Toribio:2011arXiv1103.0990T}, for nearby galaxies dwarf galaxies in the local volume \\citep{Lee:2009}.  A recent study by \\cite{Devereux:2009} adopted the Rauzy method to determine the completeness of multi-wavelength selected data.  {\\CI} extended the Rauzy test to the more realistic case of data characterised by both a faint {\\it and\\/} bright apparent magnitude limit.  Moreover we introduced a new variant statistic, denoted {\\tv}, that samples the cumulative distance modulus, $Z$, distribution but retains similar robust properties to those of {\\tc} -- i.e. being independent of the spatial distribution of galaxies.  By sampling the data in this way, the {\\tv} statistic can be considered as an improved, differential version of the the widely used {\\vmax} test (which assumes spatial homogeneity).  Our next paper, {\\CII}, highlighted the fact that the previously defined completeness estimators may be susceptible to `shot-noise' effects if studying completeness in a small parent data-set.  This phenomenon can be explained in the context of {\\CI}, where the basic construction of the  {\\tctv} statistics allocated a volume limited subsample to each individual galaxy in the catalogue. With the introduction of a secondary bright apparent magnitude limit (as illustrated in \\F{Fig:0MZ}) the size of the volume limited subsample constructed for each galaxy may be significantly restricted. {\\CII} showed how this may render the \\tctv\\,completeness estimators `shot-noise' dominated when the volume limited regions are very sparsely sampled.  As a result {\\CII} proposed an extension of the basic \\tctv completeness estimators facilitating their construction in a manner that essentially maintains a constant `signal-to-noise' ({\\sn}) level.  This article examines more deeply characteristic forms of incompleteness and how they will be manifested in the resulting {\\tctv} estimators.  We also  further extend the tools developed in {\\CII}, where we now introduce an error propagation analysis for the {\\tctv} statistics. Hence we continue our task of providing a statistically rigorous, but practical, foundation for testing the magnitude completeness of current and future redshift surveys.  The structure of this paper is as follows.  In \\S~\\ref{sec:statsframework} we briefly outline the basic framework that has underpinned the development of the {\\tctv} statistics, from R01 through to the the signal-to-noise approach detailed in {\\CII}. In \\S~\\ref{sec:character} we explore various forms of incompleteness   through the use of real and simulated data. In \\S~\\ref{sec:errorprop} we derive expressions for the error propagation of {\\tctv}.  Finally in \\S~\\ref{sec:conclusions} we discuss our findings and summarise our conclusions. ",
        "conclusions": "\\label{sec:conclusions} In the first part of this article we examined more closely the conditions under which systematic effects in survey data will manifest themselves in our completeness statistics.  To demonstrate this effectively, we firstly probed the conditions under which our completeness estimators remained robust to changes in the data.  Using actual survey data we uniformly randomly removed objects from the catalogue via a bootstrapping approach.  In a second test, we removed slices of data in either $Z$ or $M$.  This too resulted in no significant change in  both {\\tctv}. Therefore, in both cases where separability between $M$ and $Z$ was retained, {\\tctv} were shown to be robust.  Systematic effects in $m$ were then explored for  three cases which can be summarised as follows: \\begin{enumerate} \\item {\\it Under-completeness}:  This may arise when objects have not been observed over a particular redshift range and within apparent magnitude bins.  Through the use of real and mock data, we demonstrated that if this effect is present at even the level of a few percent, a characteristic signature will be observed in {\\tctv}, manifest as a significant drop in both estimators for trial apparent magnitude limits that lie within the incomplete region(s), followed by a steep rise in {\\tctv} for fainter {\\mstar} values.\\\\ \\item {\\it Over-completeness}: Here there may be more objects observed in a particular apparent magnitude range relative to the underlying distribution. In this case {\\tctv} are observed to display the opposite characteristics to those of case (i). Thus, as {\\mstar} moves across the incomplete region on the {\\mz} plane, a distinct peak in both {\\tctv} occurs followed by a distinct dip as {\\mstar} moves back into the `complete' region.\\\\ \\item {\\it Evolution}: For the particular case of luminosity evolution studied here, we demonstrated that mock magnitudes which are  drawn from an evolving LF will show a characteristic and sustained drop in {\\tc} below $-3 \\sigma$  at an {\\mstar} value significantly brighter than the actual apparent magnitude limit of the sample.  Applying the correct evolutionary model to both $M$ and $Z$, corrects this form of incompleteness. This characteristic behaviour for {\\tc} was confirmed with the use of real QSO survey data from C04 which is known to display strong evolutionary properties. \\end{enumerate} We note, however, that in cases (i) and (ii) we used a highly artificial scheme to introduce systematics to the data. In reality, one might expect some combination of these incompleteness effects to be present in the data, which may therefore partially cancel out. Nevertheless one might adopt, for example, an iterative weighting scheme to correct for any residual signatures in {\\tc} or {\\tv}.  This approach will be investigated in a later paper in the context of estimating the galaxy luminosity function.  In the first part of this paper we also revisited the impact on {\\tc} or {\\tv} of surveys that have an unmodelled secondary bright apparent magnitude limit.   The purpose of this was also to revisit our initial completeness results for the 2dF survey which led to the extension of {\\tc} for the bright limit case. We showed that imposing a successively fainter, but unmodelled, bright magnitude limit resulted in a systematic downward `shifting' in the value of the completeness statistics.  As the bright limit became fainter, the systematic shifting became more negative. We also showed that for the mock catalogues with a small systematic perturbation added to the apparent magnitude distribution, the systematic dip in {\\tc} corresponding to the affected region also remained and was, if anything, more significant.  Thus, reconsidering our 2dF results from {\\it Completness I}, we surmise that the systematic dip observed in the {\\tctv} statistics was more likely the result of some other underlying systematic effect not related to the bright limit of the survey, but which was then subsequently masked by our choice of narrow width in $\\delta Z$ and $\\delta M$,  as explored in {\\CII}.  In future work we will also explore how the properties of our random variables $\\zeta$  and $\\tau$ can be exploited to constrain luminosity evolution in galaxy surveys. In essence, $\\zeta$ and $\\tau$ are powerful probes for identifying residual correlations in $M$ and $Z$ due to evolution. Thus, to constrain evolutionary models, we can extend our methodology to include e.g.  the  Kullback-Leibler divergence \\citep{Kullback:1951} relative entropy method, which measures the difference between two probability distributions $p$ and $q$, where $p$ represents the observed distribution of the data ($\\zeta$ or $\\tau$ in our case) and $q$ represents our theoretical model for that distribution (i.e. after application of a specific correction for luminosity evolution).  In the final part of this paper we presented a full error propagation and covariance analysis for the {\\tctv} statistics, further developing the ``adaptive smoothing\" procedure introduced in {\\CII}.  By addressing all these issues we believe that we have laid  a very comprehensive foundation for testing magnitude completeness limits that will become crucial for the next generation of redshift surveys."
    },
    "1107/1107.3406_arXiv.txt": {
        "abstract": "{Cosmic ray antiprotons provide an important probe for the study of cosmic-ray propagation in the interstellar space and to investigate the existence of Galactic dark matter.  The ARGO-YBJ experiment, located at the Yangbajing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm$^2$ ), is the only experiment exploiting the full coverage approach at very high altitude presently at work. The ARGO-YBJ experiment is particularly effective in measuring the cosmic ray antimatter content via the observation of the cosmic rays Moon shadowing effect. Based on all the data recorded during the period from July 2006 through November 2009 and a full Monte Carlo simulation, we searched for the existence of the shadow produced by antiprotons at the few-TeV energy region. No evidence of the existence of antiprotons was found in this energy region. Upper limits to the antip/p flux ratio are set to 5\\% at a median energy of 2 TeV and 6\\% at 5 TeV with a confidence level of 90\\%. In the few-TeV energy range this result is the lowest available.} ",
        "introduction": "Very High Energy Cosmic Ray (VHE CR) antiprotons are an essential diagnostic tool to approach the solution of several big topics of cosmology, astrophysics and particle physics, besides for studying fundamental properties of the CR sources and propagation medium \\cite{stecker02}. The enigma of the matter/anti-matter asymmetry in the local Universe, namely the existence of antimatter regions, the signatures of physics beyond the standard model of particles and fields, as well as the determination of the essential features of CR propagation in the insterstellar medium, are a few research topics which would greatly benefit from the detection of VHE antiprotons \\cite{stecker84,klopov00,strong07,pamelap10,cirelli,donato09,golden}.  Cosmic rays are hampered by the Moon, therefore a deficit of CRs in its direction is expected (the so-called \\emph{Moon shadow}). Moreover, the Earth-Moon system acts as a magnetic spectrometer. In fact, due to the Geomagnetic field (GMF) the Moon shadow shifts westward by an amount depending on the primary CR energy. The paths of primary antiprotons are therefore deflected in the opposite sense in their way to the Earth. This effect allows, in principle, the search for antiparticles in the opposite direction of the observed Moon shadow.  If the energy is low enough and the angular resolution is good we can distinguish the two shadows, one shifted towards West due to the protons and the other shifted towards East due to the antiprotons. At high energies ($\\geq$ 10 TeV) the magnetic deflection is too small compared to the angular resolution and the two shadows cannot be disentangled. At low energies ($\\leq$500 GeV) the shadows are much deflected but washed out by the poor angular resolution, thus the sensitivity is limited. Therefore, there is an optimal energy window for the measurement of the antiproton abundance.  The ARGO-YBJ experiment is especially recommended for the measurement of the CR antimatter content via the observation of the Galactic CR shadowing effect due to: (1) good angular resolution and pointing accuracy and their long term stability; (2) low energy threshold; (3) real sensitivity to the GMF. Indeed, the low energy threshold of the detector allows the observation of the shadowing effect with a sensitivity of about 9 standard deviations (s.d.) per month at TeV energy.  In this paper we report the measurement of the CR $\\overline{p}/p$ flux ratio in the TeV energy region with all the data collected during the period from July 2006 to November 2009. ",
        "conclusions": "The ARGO-YBJ experiment is observing the Moon shadow with high statistical significance at an energy threshold of a few hundred GeV. Using all data collected until November 2009, we set two upper limits on the $\\bar{p}/p$ flux ratio: 5\\% at an energy of 1.4 TeV and 6\\% at 5 TeV with a confidence level of 90\\%. In the few-TeV range the ARGO-YBJ results are the lowest available, useful to constrain models for antiproton production in antimatter domains."
    },
    "1107/1107.6015_arXiv.txt": {
        "abstract": "We investigate NLTE effects for magnesium and calcium in the atmospheres of late-type giant and supergiant stars. The aim of this paper is to provide a grid of NLTE/LTE equivalent width ratios $W/W^*$ of Mg and Ca lines for the following range of stellar parameters: $T_{\\rm {eff}} \\in [3500, 5250]$~K, $\\log\\ g \\in [0.5, 2.0]$~dex and [Fe/H$] \\in [-4.0, 0.5]$~dex. We use realistic model atoms with the best physics available and taking into account the fine structure. The Mg and Ca lines of interest are in optical and near IR ranges. A special interest concerns the lines in the {\\it Gaia} spectrograph (RVS) wavelength domain [8470, 8740]~\\AA. The NLTE corrections are provided as function of stellar parameters in an electronic table as well as in a polynomial form for the {\\it Gaia}/RVS lines. ",
        "introduction": "Individual stars of different ages carry information that allow astronomers to detail the sequence of events which built up the Galaxy. Calcium (Ca) and magnesium (Mg) are key elements for the understanding of these events since these $\\alpha$-elements are mainly produced by supernovae of type II. Therefore, they are both important tracers to derive trends of chemical evolution of the Galaxy. These two elements are easily observed in spectra of late-type stars through strong or weak lines. Among them, there are the \\mg\\ b green triplet, the \\caI\\ red triplet and the \\caII\\ IR triplet (CaT) lines which are measurable in the most extreme metal-poor stars \\citep{Christlieb02} even at moderate spectral resolutions.  The lines of these $\\alpha$-elements provide diagnostic tools to probe a wide range of astrophysical mechanisms in stellar populations. The study of $\\alpha$-element synthesis in halo metal-poor stars constrains the rate of SNe II \\citep{Nakamura99}. The Mg element is a reliable reference of the early evolutionary time scale of the Galaxy (e.g. \\citealt{Cayrel04}, \\citealt{Gehren06} and \\citealt{Andrievsky10}). Moreover, the \\caI\\ triplet at $\\lambda \\lambda$ 6102, 6122 and 6162, can be used as a gravity indicator in detailed analyses of dwarf and subgiant stars \\citep{Cayrel96}. Two Ca lines of different ionization stages can also be used to constraint surface gravity in extremely metal-poor stars (\\caI\\ 4226~\\AA\\ and \\caII\\ 8498~\\AA) as shown by \\citet{Mashonkina07}. Moreover, the line strengths of CaT lines can be used as metallicity indicator for stellar population in distant galaxies \\citep{Id97} or to search for the most extreme metal-poor stars in the vicinity of nearby dwarf spheroidal galaxies \\citep{Starkenburg10}. The CaT lines are also used to probe stellar activity by the measure of the central line depression as done by \\citet{Andretta05} and \\citet{Busa07}.  It is well known that the assumption of Local Thermodynamic Equilibrium (LTE) is the main source of uncertainties (after the quality of the atomic data). In his review, Asplund (2005) summarized the need to take Non-LTE (NLTE) effects into account to get reliable and precise chemical abundances. Such effects are crucial in the atmospheres of metal-poor stars due to several mechanisms: the overionization in the presence of a weak opacity, the photon pumping and the photon suction. However, since the NLTE computations are still time consuming as soon as we deal with realistic  model atoms,  most of the studies are still done under the LTE assumption for large surveys (e.g., \\citealt{Edvardsson93} for late-type dwarf stars, \\citealt{Liu10} and \\citealt{DeSilva11} for late-type giant and subgiant stars).  Several studies of NLTE effects on \\mg\\ lines were performed for stars of various metallicities \\citep{Mashonkina96, Zhao98, Zhao00, Gratton99, Shimanskaya00, Idiart00, Mishenina04, Andrievsky10}. All studies showed an increase of NLTE effects with decreasing metallicity which imply corrections up to $+0.2$~dex. The first attempt to build a realistic model atom of \\mg\\ was performed by \\citet{Carlsson92} to explain the \\mg\\ solar emission lines near to 12 $\\mu$m. They demonstrated the photospheric nature of these lines by using the photon suction mechanism which required a model atom with Rydberg states to ensure the collisional-dominated coupling with the ionization stage. The most complete Mg model atom was built by \\citet{Przybilla01} with 88 energy levels for \\mg\\ and 37 levels for \\mgII.  However, they did not take into account the fine structure of the triplet system. The first \\mg\\ model atom that includes the fine structure was built by \\citet{Idiart00} who applied it to the determination of the Mg abundance for more than two hundred late-type stars. We can mention an update of the \\mg\\ model atom of \\citet{Carlsson92} done by \\citet{Sundqvist08} in order to reproduce IR emission lines in K giants, taking into account the electronic collisions between states with the principal quantum number $n \\le 15$.  NLTE effects on Ca were investigated by \\citet{Jorgensen92} who gave a relation between equivalent widths of CaT lines as functions of the surface gravities and the effective temperatures. As for the \\mg, \\citet{Idiart00} included the fine structure in their model atom of \\caI. They combined these results to the NLTE Fe abundance predictions \\citep{Thevenin99} and performed an NLTE analysis of the same set of late-type dwarf and subgiant stars. They concluded that the [Ca/Fe] remained mostly unchanged but some parallel structures appeared in the NLTE [Ca/Fe] vs [Fe/H] diagram. The most complete Ca atom was done by \\citet{Mashonkina07} who built a model atom through two ionization stages with 63 energy levels for \\caI\\ and 37 levels for \\caII, including fine structure only for the first states. They used the most recent and accurate atomic data from the Opacity Project \\citep{Seaton94}. They concluded that NLTE abundance corrections can be as large as 0.5 dex (for the \\caI\\ 4226 \\AA\\ resonance line) in extremely metal-poor stars. Using these NLTE corrections, \\citet{Asplund09} found good agreement of solar \\caI\\ and \\caII\\ abundances.  In the view of the ESA {\\it Gaia} mission (\\citealt{Wilkinson05}), it is important to investigate such NLTE abundance corrections in order to distinguish different stellar populations. The onboard spectrograph, named RVS (Radial Velocity Spectrometer, \\citealt{Katz04}) which was designed primarily to get the radial velocities can also be used in principle to extract individual chemical abundances. The spectral window considered was selected in the near-infrared [8470-8740]~\\AA\\ because of the presence of the CaT lines which will be visible even in the most metal-poor stars and at a considered low resolution ($R~\\sim$~11500 at best). The counterpart ground-based observations is the RAVE survey \\citep{Steinmetz06} which will provide about 1 million of spectra. These two surveys will analyze a representative sample of stars throughout the Galaxy, allowing an unprecedented kinematic and chemical study of the stellar populations. The huge amount of stars will prevent us to analyze them one by one. It is necessary to have precalculated correction tables that can be used for automatic stellar parameter determination and this is the aim of this work.   The structure of the paper is as follows: in Sect.~2, we  present the model atoms of \\mg, \\caI\\ and \\caII. In Sect.~3, we test these models by comparing NLTE synthetic lines with several solar (Sect.~3.1) and Arcturus (Sect.~3.2) lines of \\mg\\ and \\caII\\ in the {\\it Gaia}/RVS region. Then, in Sect.~4.1, we present a grid of NLTE corrections of equivalent widths of spectral lines, selected in Sect.~4.2, useful for the {\\it Gaia} mission and the RAVE survey. The behaviours of \\mg, \\caI\\ and \\caII\\ NLTE lines, the comparisons with previous works, the estimate of uncertainties on atomic data, and the polynomial fits of the {\\it Gaia}/RVS lines are discussed in Sect.~4. ",
        "conclusions": "We have performed NLTE computations for the \\mg, \\caI\\ and \\caII\\ model atoms in late-type giant and supergiant stars. We provide NLTE/LTE EW ratios for a grid of 453 \\marcs\\ model atmospheres ($3500\\le T_{{\\rm eff}}\\le 5250$~K, $0.5\\le \\log g \\le 2.0$~dex and $-4.00\\le [$Fe/H$] \\le+0.50$~dex) in electronic forms\\footnote{available on the Vizier/CDS database}. The model atoms are based on the assumption that we do not take into account inelastic collisions with neutral hydrogen since realistic quantum mechanical collisional cross-sections are still unavailable. We used a formulation for electronic collisions that underestimate the collisional rate. % Such underestimation induced an upper limit on the NTLE/LTE EW ratios especially for the \\mg\\ lines in the {\\it Gaia}/RVS wavelength range. The use of fine structure in the model atoms do not affect strongly NLTE results since the $W/W^*$ of the components of a multiplet are very similar but permit a consistent representation of the physics. This work will be extended to late-type main sequence stars when collisional cross-sections with hydrogen are available.  The main conclusions for the lines outside the {\\it Gaia}/RVS wavelength range are as follows. For the \\mg\\ and \\caI\\ lines, the assumption of LTE underestimates the Mg and the Ca abundances and gives mainly positive NLTE abundance correction. For the \\caII\\ lines, the assumption of LTE overestimates the Ca abundance and gives mainly negative NLTE abundance corrections. However, most of the \\mg\\ lines show $W/W^* > 1$ that can reach 2 for the coolest models with an increase of the sensitivity to the surface gravities. The \\mg\\ b and \\caI\\ red triplets show $W/W^* < 1$ that can reach 0.5 for the most metal-poor and the hottest model atmospheres. The NLTE effects for the \\mg~4571~\\AA\\ and the \\caI~6572~\\AA\\ intercombination and resonance lines are very sensitive to the metallicity and to the surface gravity. The NLTE effects on \\caII\\ H\\&K lines are negligible except at [Fe/H$] = -4$.  For the {\\it Gaia}/RVS lines, NLTE computations give the following trends. The very weak 8473~\\AA\\ line is mainly formed in LTE whereas the \\mg\\ 8736~\\AA\\ and the \\mg\\ $\\lambda\\lambda$ 8710, 8712 and 8717 triplet lines are mainly formed in NLTE with a strong sensitivity to the surface gravity for the triplet. The weak \\caI\\ triplet lines are mainly formed in NLTE but vanish as soon as [Fe/H$] < 1$. The famous CaT lines are mainly formed in LTE if [Fe/H$] \\ge -2$. The NLTE effects increase with a decrease of the metallicity and with an increase of the surface gravity. We show that the $W/W^*$(CaT) can increase by 20\\% for $\\log g$ varying from 0.5 to 2.0. For convenience, we provide a polynomial fit computed for the {\\it Gaia}/RVS lines: the \\mg\\ $\\lambda\\lambda$ 8473, $\\left<8715\\right>$ triplet, 8736 multiplet, the \\caI\\ $\\lambda\\lambda$ 8525, 8583 and 8633, and the CaT lines. The polynomial can be extensively used by the automatic tools dedicated to the analysis and extract chemical abundances for millions of stars in the context of the large surveys as, for example, {\\it Gaia} and RAVE."
    },
    "1107/1107.2827_arXiv.txt": {
        "abstract": "We review an approach to observation-theory comparisons we call ``Taste-Testing.\"  In this approach, synthetic observations are made of numerical simulations, and then both real and synthetic observations are ``tasted\"  (compared) using a variety of statistical tests.   We first lay out arguments for bringing theory to observational space rather than observations to theory space.  Next, we explain that generating synthetic observations is only a step along the way to the quantitative, statistical, taste tests that offer the most insight.   We offer a set of examples focused on polarimetry, scattering and emission by dust, and spectral-line mapping in star-forming regions.   We conclude with a discussion of the connection between statistical tests used to date and the physics we seek to understand.  In particular, we suggest that the ``lognormal\" nature of molecular clouds can be created by the interaction of many random processes, as can the lognormal nature of the IMF, so that the fact that both the ``Clump Mass Function\" (CMF) and IMF appear lognormal does not necessarily imply a direct relationship between them. ",
        "introduction": "Saying that one is working on ``comparing observations with simulations\" sounds straightforward.  Alas, though, comparing astronomical experimental data (``observations\") about star formation with relevant theory is much more difficult than [experiment]:[theory] comparisons in almost any other area of physics. The reasons are twofold. First, the theories involved are not typically ``clean\" analytic ones, but instead require simulation to make predictions. And second, the measurements observers offer involve complicated-to-interpret ``fluxes\" from various physical processes rather than more direct measures of physical parameters.  Let's say a theory makes a prediction about density.  In a laboratory, one fills a container of known interior volume and weight with the substance under study using some kind of simple mechanical device (e.g. a spoon), puts the container on a well-calibrated scale, measures a weight, subtracts the weight of the container, divides by the well-known gravitational constant and the known volume, and {\\it voil\\`a}, density.   In a star-forming region...not so simple.  One points a telescope of (sometimes poorly) known beam response at some patch of 2D sky and collects all the photons that come to a detector of (sometimes poorly known or variable) efficiency via the telescope during a well-measured period of time.  Such an observation gives the ``flux\" (or lack thereof) of some quantity (e.g. thermal emission from dust, emission from a spectral line in gas, dust extinction).  Next, that flux needs to be converted either to a 2D ``column density\" on the sky, or a 3D volume density (only possible using chemical excitation models for spectral lines).  In the case of a column density, the ``column\" needs to be assigned a length to extract a volume density.  And, in either case, it is not clear what mixture of actual densities conspire along a line of sight to give the measured ``representative\" volume density.  So, should we give up on observational measurements of physical parameters in star formation research?  Of course not.  But, we should consider that perhaps comparisons of observations and theory (primarily simulations) are best carried out in an ``observational space,\" where synthetic observations of theoretical output are made in order to enable more direct statistical comparison with ``real\" (observational) data.   This comparison in observational space has been referred to in the past by some (including me) as ``Taste-Testing.\"  If one's goal is to reproduce the recipe for a great soup eaten in a restaurant, then a good course of action is to guess the ingredients based on dining experience (observation) and the processes based on cooking experience (physics) create the soup, take a look at it, and then if it seems to look and smell right, then finally ``taste\" it.  Centrifuging the soup--as a theorist might try to do with observations in ``theorists'\" space--to discern its ingredients might give you a basic chemical breakdown, but it would not tell you how to make it again.   Thus, we define a legitimate [observation]:[theory] taste test as one where statistical (taste) tests are applied in observer's (edible) space.  An overview of the process is shown in Fig.\\,\\ref{fig1}.  \\begin{figure}[h] \\includegraphics[width=5.5in]{figure_tests.pdf} \\vspace*{0.2 cm} \\caption{The Taste-Testing Process.  The processes leading to measurements are represented by the green arrow for ``real\" data and by the two purple arrows for synthetic data.  The box at right shows a comparison of the fraction of self-gravitating (according to a virial analysis) material as a function of scale in the simulations of \\citet{Padoan:2006p711} and in the L1448 region of Perseus (see \\citet {Rosolowsky:2008p53, GoodmanNat09} and \\S 3, below.) } \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "1107/1107.2685_arXiv.txt": {
        "abstract": "We investigate differential rotation in rapidly rotating solar-type stars by means of an axisymmetric mean field model that was previously applied to the sun. This allows us to calculate the latitudinal entropy gradient with a reasonable physical basis. Our conclusions are as follows: (1) Differential rotation approaches the Taylor-Proudman state when stellar rotation is faster than solar rotation. (2) Entropy gradient generated by the attached subadiabatic layer beneath the convection zone becomes relatively small with a large stellar angular velocity. (3) Turbulent viscosity and turbulent angular momentum transport determine the spatial difference of angular velocity $\\Delta\\Omega$. (4) The results of our mean field model can explain observations of stellar differential rotation. ",
        "introduction": "Our sun has an eleven-year magnetic activity cycle, which is thought to be sustained by the dynamo motion of internal ionized plasma, i.e., a transformation of kinetic energy to magnetic energy \\citep{1955ApJ...122..293P}. Our understanding of the solar dynamo has significantly improved during the past fifty years, and some kinematic studies can now reproduce solar magnetic features such as equatorward migration of sunspots and poleward migration of the magnetic field \\citep{1995A&A...303L..29C,1999ApJ...518..508D,2001A&A...374..301K,2005LRSP....2....2C,2010ApJ...709.1009H,2010ApJ...714L.308H}. The most important mechanism of the solar dynamo is the $\\Omega$ effect, the bending of pre-existing poloidal magnetic fields by differential rotation and the generation of toroidal magnetic fields. Thus, the distribution of the differential rotation in the convection zone is a significant factor for the solar dynamo. Using helioseismology, it has recently been shown that the solar internal differential rotation is in a non-Taylor-Proudman state \\citep[see review by][]{2003ARA&A..41..599T}, meaning the iso-rotation surfaces are {\\it not} parallel to the axis. \\par Based on solar observations, it is known that Ca H-K fluxes can be a signature of stellar chromospheric activity, and such chromospheric signatures are in correlation with magnetic activity. \\cite{1968ApJ...153..221W,1978ApJ...226..379W} and \\cite{1995ApJ...438..269B} discuss a class of stars that shows a periodic variation in Ca H-K fluxes, which suggests that they have a magnetic cycle similar to our sun. It is natural to conjecture that such magnetic activity is maintained by dynamo action. Various studies have been conducted to investigate the relationship between stellar angular velocity $\\Omega_0$ and its latitudinal difference $\\Delta\\Omega$ i.e., $\\Delta\\Omega\\propto\\Omega_0^n$, where the suggested range of $n$ is $0 <n < 1$ \\citep{1996ApJ...466..384D,2003A&A...398..647R,2005MNRAS.357L...1B}. This means that the angular velocity difference $\\Delta\\Omega$ increases and the relative difference $\\Delta\\Omega/\\Omega_0$ decreases with increases in the stellar rotation rate $\\Omega_0$.\\par In this paper, we investigate differential rotation in rapidly rotating stars using a mean field framework. Our study is based on the work of \\cite{2005ApJ...622.1320R}, in which he suggests the importance of the role of the subadiabatic layer below the convection zone in order to maintain a non-Taylor-Proudman state in the Sun. The aim of this paper is to use a mean field  model to analyze firstly the dependence of the morphology of differential rotation on stellar angular velocity, and secondly the physical process which determines the observable angular velocity difference $\\Delta \\Omega$. According to our knowledge, this is the first work which systematically discusses the application of Rempel's (2005b) solar model to stars. \\par Other research adopts another approach to the use of mean field models for the analysis of differential rotation in rapidly-rotating stars \\citep{1995A&A...299..446K,2001A&A...366..668K}. In these studies, the non-Taylor-Proudman state is sustained by anisotropy of turbulent thermal conduction. This anisotropy is generated by the effects of stellar rotation on convective turbulence.\\par Three-dimensional numerical studies on stellar differential rotation also exist \\citep{2008ApJ...689.1354B,2009AnRFM..41..317M}. In these studies, they resolve stellar thermal driven convection and can calculate a self-consistent turbulent angular momentum transport and anisotropy of turbulent thermal conductivity. The subadiabatic layer below the convection zone, however, is not included. The effects of anisotropy of turbulent thermal conductivity and the subadiabatic layer are discussed in this paper. ",
        "conclusions": "We run simulations for seventeen cases, with Table \\ref{param} showing the parameters for each case. \\subsection{Stellar Differential Rotation}\\label{taylor} In this section, we discuss the cases with angular velocities up to 16 times the solar value (represented by $\\Omega_\\odot$), placing an emphasis on the morphology of stellar differential rotation. Fig. \\ref{rapid} shows the results of our calculations which correspond to cases 1-5 in Table \\ref{param}. It is found that the larger stellar angular velocity is, the more likely it is for differential rotation to be in the Taylor-Proudman state, in which the contour lines of the angular velocity are parallel to the rotational axis. To evaluate these results quantitatively, we define a parameter which denotes the morphology of differential rotation. We call it the Non-Taylor-Proudman parameter (hereafter the NTP parameter), which is expressed as \\begin{eqnarray} P_\\mathrm{ntp}=\\frac{1}{R_\\odot^2\\Omega_0^2}\\int\\frac{\\partial \\Omega_1^2}{\\partial z}dV =\\frac{1}{R_\\odot^2\\Omega_0^2}\\int \\left( \\cos\\theta\\frac{\\partial }{\\partial r} -\\frac{\\sin\\theta}{r}\\frac{\\partial }{\\partial \\theta} \\right) \\Omega_1^2dV, \\end{eqnarray} where $\\Omega_0$ is the angular velocity of the radiative zone. When the NTP parameter is zero, differential rotation is in the Taylor-Proudman state. Conversely, differential rotation is far from the Taylor-Proudman state with a large absolute value of the NTP parameter. The value of the NTP parameter with various stellar angular velocities is shown in Fig. \\ref{npp}. The NTP monotonically decreases with increases in stellar angular velocity. These results indicate that with large stellar angular velocity values, differential rotation approaches the Taylor-Proudman state. These results are counter-intuitive, however, since we do not expect differential rotation to approach the Taylor-Proudman state with increasing stellar angular velocity values, since the $\\Lambda$ effect, which is a driver of the deviation from the Taylor-Proudman state, is proportional to stellar angular velocity $\\Omega_0$. These are the most significant findings of this paper, so hereafter in this section we discuss these unexpected results. \\par We next discuss the temperature difference between the equator and the pole at the base of the convection zone ($r=0.71R_\\odot$). Since temperature is given as a function of entropy by \\begin{eqnarray} T_1=\\frac{T_0}{\\gamma} \\left[s_1+(\\gamma-1)\\frac{p_1}{p_0} \\right], \\end{eqnarray} and it is easier to measure than entropy, we use it here for discussing the thermal structure of the simulation results in the convection zone. Further, although it is mentioned in \\S \\ref{differential} that entropy gradient is crucial for breaking the Taylor-Proudman constraint, the temperature difference can be used as its proxy. Fig. \\ref{entropy} shows the relationship between stellar angular velocity $\\Omega_0$ and temperature difference $\\Delta T$ at $r=0.71R_\\odot$, where $\\Delta T =\\max (T_1(r_\\mathrm{bc},\\theta))-\\min (T_1(r_\\mathrm{bc},\\theta))$. Although the temperature difference monotonously increases with larger stellar angular velocity values, it is not enough to make the rotational profile largely deviate from the Taylor-Proudman state. This can be explained by using the thermal wind equation, which is a steady state solution of eq. (\\ref{therm01}): \\begin{eqnarray} 0= r\\sin\\theta\\frac{\\partial \\Omega^2}{\\partial z} -\\frac{g}{\\gamma r}\\frac{\\partial s_1}{\\partial \\theta}.\\label{therm02} \\end{eqnarray} The inertial term and the diffusion term are neglected here. This equation indicates that, for a given value of the NTP, we need an entropy gradient proportional to $\\Omega_0^2$. However, our simulation results show that $\\Delta T \\propto \\Omega_0^{0.58}$, which means that as $\\Omega_0$ increases, the thermal driving force becomes insufficient to push differential rotation away from the Taylor-Proudman state. In other words, the latitudinal entropy gradient in rapidly rotating stars is so small that differential rotation stays close to the Taylor-Proudman state. In our model, meridional flow generates latitudinal entropy gradient at the base of the convection zone. It is conjectured that the insufficient thermal drive is due to a slow meridional flow.\\par We next investigate the dependence of meridional flow on stellar angular velocity. Fig. \\ref{vari_some} shows the radial profile of latitudinal velocity $v_\\theta$ at $\\theta=45^\\circ$, using the results of cases 1, 2 and 9. In case 2, stellar angular velocity is twice that of case 1 (the solar value). In case 9, stellar angular velocity is equal to the solar value, and the amplitude of the $\\Lambda$ effect is two times the value in case 1. Fig. \\ref{vari_some} shows that meridional flow does not depend on stellar angular velocity, while it correlates with the $\\Lambda$ effect. Considering eq. (\\ref{lambda}), the $\\Lambda$ effect increases with larger values of stellar angular velocity, since the amplitude of the $\\Lambda$ effect is proportional to $\\Omega_0$. The reason why differential rotation in rapidly rotation stars is close to  the Taylor-Proudman state is that meridional flow does not become fast with large stellar angular velocity values. \\par We interpret the result that the speed of meridional flow does not depend on stellar angular velocity in our model as follows. With large values of stellar angular velocity, more angular momentum is transported by the $\\Lambda$-effect (Note that the $\\Lambda$-effect is proportional to $\\Omega_0$ in equation (\\ref{amp_lambda})), so meridional flow obtains more energy from differential rotation. The energy gain does not result in an increase in speed because of the associated enhancement of the Coriolis force, which bends the meridional flow in the longitudinal direction. Another explanation is possible in terms of angular momentum transport. The angular momentum fluxes from both meridional flow and the Reynolds stress ($\\Lambda$ effect) must be balanced in a steady state. The former is proportional to $v_\\mathrm{m}\\Omega_0$ and the latter is proportional to $\\Omega_0$, where $v_\\mathrm{m}$ is the amplitude of meridional flow. Therefore, meridional flow does not depend on stellar angular velocity \\citep{2005LRSP....2....1M}. Our results (Fig. \\ref{vari_some}) indicate that with a larger stellar angular velocity (case 2), the above mechanism does not generate fast meridional flow. However, this does not occur  when only the $\\Lambda$ effect is large (case 9). \\subsection{Angular Velocity Difference on the Surface}\\label{delta_omega} In this subsection we discuss angular velocity difference $\\Delta \\Omega$ at the surface and the relationship between our results and previous observations. We conduct numerical simulations to investigate the physical process which determines $\\Delta \\Omega$ (cases 1, 6-11). We define angular velocity difference as $\\Delta \\Omega = \\max(\\Omega_1(r_\\mathrm{max},\\theta))-\\min(\\Omega_1(r_\\mathrm{max},\\theta))$. \\par $\\Delta \\Omega$ is determined by two opposing effects, a smoothing effect from turbulent viscosity and a steepening effect from the $\\Lambda$ effect. In a stationary state these two effects cancel each other out. Latitudinal flux for turbulent viscosity and the $\\Lambda$ effect can be written as $\\rho_0\\nu_\\mathrm{0v}\\Delta \\Omega/\\Delta \\theta$ and $\\rho_0\\nu_\\mathrm{0l}\\Lambda_0\\Omega_0$, respectively. Because these two have approximately the same value, $\\Delta\\Omega$ can be estimated as \\begin{eqnarray} \\Delta \\Omega \\sim\\frac{\\nu_\\mathrm{0l}}{\\nu_\\mathrm{0v}}\\Lambda_0\\Omega_0\\Delta\\theta, \\label{deltaomega} \\end{eqnarray} where $\\Delta \\theta$ denotes the differential rotation region.\\par In order to confirm eq. (\\ref{deltaomega}), we conduct two sets of simulations, firstly varying the value of turbulent viscosity ($\\nu_\\mathrm{0v}$), and secondly the amplitude of the $\\Lambda$ effect ($\\Lambda_0$). Note that the setting for turbulent viscosity does not reflect a real situation, since the coefficients of turbulent viscosity and the $\\Lambda$ effect should have a common value. Nonetheless, this is necessary for the purpose of our investigation. The simulation results are shown in Figures \\ref{vis_vari} and \\ref{lam_vari}. We obtain $\\Delta \\Omega \\propto \\nu_\\mathrm{0v}^{-0.88}$ and $\\Delta \\Omega \\propto \\Lambda_0^{1.1}$, which are consistent with eq. (\\ref{deltaomega}). \\par Fig. \\ref{angular_difference} shows the results of the dependence of $\\Delta \\Omega$ on $\\Omega_0$ (Cases 1-5). Asterisks denote the difference at the surface between the equator and the pole, squares show the difference between the equator and the colatitude $\\theta=45^\\circ$, and triangles are the difference between the equator and the colatitude $\\theta=60^\\circ$. The difference at low latitudes  (squares and triangles) monotonically increases with stellar angular velocity. However this is not the case for angular velocity difference between the equator and the pole (asterisk). As we discussed in \\S \\ref{taylor}, when stellar rotation velocity is large, the Taylor-Proudman state is achieved, meaning the gradient of angular velocity at the surface concentrates in lower latitudes. Due to this concentration, $\\Delta\\theta$ becomes smaller in Eq. (\\ref{deltaomega}) with larger values of $\\Omega_0$. Thus, $\\Delta \\Omega_0$ does not show an explicit dependence on $\\Omega_0$. At low latitudes, $\\Delta\\theta$ is fixed and the angular velocity difference increases with stellar angular velocity. We obtain $\\Delta\\Omega\\propto\\Omega_0^{0.43}$ (between the equator and the colatitude $\\theta=45^\\circ$: squares) and $\\Delta\\Omega\\propto\\Omega_0^{0.55}$ (between the equator and the colatitude $\\theta=60^\\circ$: triangles). This indicates that $\\Delta\\Omega/\\Omega_0 $ decreases with stellar angular velocity. These results are consistent with previous stellar observations \\citep{1996ApJ...466..384D,2003A&A...398..647R,2005MNRAS.357L...1B}. \\subsection{Variation of $\\Lambda$-effect and superadiabaticity}\\label{variation} In this section, we discuss the dependence of meridional flow and differential rotation on free parameters. The parameter set is shown in Table \\ref{param} (cases 12-17). At first we investigate the influence of the variation of the $\\Lambda$ effect. The $\\Lambda$ effect has two free parameters, i.e., amplitude $\\Lambda_0$ and inclination angle $\\lambda$ (see \\S \\ref{s:lambda}). Amplitude is thought to become smaller with a larger stellar angular velocity, due to the saturation of the correlations such as $\\langle v'_rv'_\\phi \\rangle$ and $\\langle v'_\\theta v'_\\phi \\rangle$, where $v_r'$, $v_\\theta'$ and $v_\\phi'$ are the radial, latitudinal and longitudinal component turbulent velocities, respectively. Fig. \\ref{vari_some} shows that meridional flow becomes slower with a smaller $\\Lambda_0$, keeping the $\\Omega_0$ value constant (Case 10). It is clear with the result of \\S \\ref{taylor} that meridional flow becomes slow with a larger angular velocity when the variation of $\\Lambda_0$ is included. \\cite{2008ApJ...689.1354B} reported this effect with their three-dimensional hydrodynamic calculation. When meridional flow is slow, the entropy gradient generated by the subadiabatic layer is small, and differential rotation approaches the Taylor-Proudman state.\\par The inclination angle is thought to be small with large stellar angular velocity values, since the motion across the rotational axis is restricted \\citep{1993A&A...276...96K}. In case 12, differential rotation with a small inclination angle ($\\lambda=2.5^\\circ$) is calculated. Other parameters are the same as case 1. The radial distribution of meridional flow is shown in Fig. \\ref{vari_some}. Meridional flow becomes faster with a smaller inclination angle. Because of the efficient angular momentum transport in the $z$ direction when the inclination angle is small, the second term on the right hand side of Eq. (\\ref{therm01}) is large. This generates a large $\\omega_\\phi$, i.e. fast meridional flow. \\par In summary, we found that rapid stellar rotation causes two opposing effects on the speed of meridional flow. The speed is reduced by the suppression of $\\Lambda_0$, while it is enhanced by the angular momentum transport along the axial direction with a smaller $\\lambda$. Although the results of the three-dimensional calculation suggest that meridional flow becomes slower with a larger stellar angular velocity, our model cannot draw a conclusion about the speed of meridional flow in rapidly rotating stars. \\par Next we investigate the influence of superadiabaticity in the convection zone. In cases 13-17, superadiabaticity in the convection zone $\\delta_\\mathrm{c} = 1\\times 10^{-6}$. The differences of the NTP parameters with adiabatic and superadiabatic convection zones  $(P_{\\mathrm{ntp}(\\delta_\\mathrm{c}=0)}-P_{\\mathrm{ntp}(\\delta_\\mathrm{c}=10^{-6})})/P_{\\mathrm{ntp}(\\delta_\\mathrm{c}=0)}$ are shown in Fig. \\ref{super}. The NTP parameter values with a superadiabatic convection zone are smaller than those with an adiabatic convection zone, since meridional flow in the superadiabatic convection zone makes the entropy gradient small.  This result is suggested by \\cite{2005ApJ...631.1286R}. Note that the difference between the values of the NTP parameters with an adiabatic and those with a superadiabatic convection zone decreases as the stellar angular velocity increases, since the generation of entropy gradient by the subadiabtic layer becomes ineffective with a larger stellar angular velocity."
    },
    "1107/1107.4892_arXiv.txt": {
        "abstract": "{GRIPS is one example of next generation telescopes proposed for astronomy the energy range between hard X-ray mirror instruments such as NuStar and the Fermi telescope. The Compton telescope principle is an advantageous concept in view of background suppression, imaging sensitivity within a large field of view and energy range, and capability to measure polarization. The diversity of astrophysical sources at high energies (diffuse emission from cosmic-ray interactions, nuclear lines from point-like and diffuse sources, accreting binaries, cosmic-ray acceleration sites, novae and supernovae, GRBs) presents a challenge, and in particular emphasizes the need for large fields of view and surveys. We  discuss the astrophysical challenges which are expected to remain after the extended INTEGRAL mission, and how such a next-generation survey at low-energy gamma-rays would impact on these. We argue that qualitatively new and more direct insights could be obtained on cosmic high-energy phenomena and their underlying physical processes. }  \\FullConference{8$^{th}$ INTEGRAL Science Workshop\\\\ {\\it The Restless Gamma-Ray Universe} \\\\ October 27-30, 2010\\\\ Dublin, Ireland}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "% Astrophysics of high-energy processes needs to close the observational gap remaining from the relatively poor sensitivity obtained so far in the regime between hard X-rays and GeV or higher-energy gamma-rays. This range of the electromagnetic spectrum of cosmic photons holds unique messages which cannot be obtained otherwise, and also allows to measure directly several key processes which currently are claimed to be understood or explored yet depend on sometimes uncertain models or assumed parameters. GRIPS is one of the ideas to address this, and to invest thus into qualitatively new astronomy.  {\\bf Acknowledgements} The GRIPS consortium has discussed the astrophysical issues to be addressed and has assembled all project-related design and planning material. We appreciate the constructive work of this consortium over the past years, including the preparation of mission idea proposals for ESA's Cosmic Vision scheme."
    },
    "1107/1107.0804.txt": {
        "abstract": " ",
        "introduction": "Equation of state (EOS) for matter inside compact stars is sometimes characterized by the phase transitions. There have been expected various first-order phase transitions in nuclear matter, such as liquid-gas transition in the low densities, meson condensations and hadron-quark deconfinement transitions at high-density region \\cite{sha,tama,glen,hae}. If such phase transitions occurs inside compact stars, they affect EOS as well as thermal, dynamical or magnetic properties \\cite{bec,sag,wata09,sot}. There has been a theoretical development and recent studies have shown that the inhomogeneous states called ``pasta'' may emerge as a mixed phase during the phase transitions, based on the Gibbs conditions \\cite{revpasta}. Historically, the mixed phases have been mostly treated very naively by applying the Maxwell construction to get the EOS in thermodynamic equilibrium. Ravenhall et al.\\ discussed the possible emergence of the pasta structures at low-density nuclear matter, since uniform nuclear matter is not stable at low densities \\cite{Rav83}. Their pasta phase may be regarded as a mixed phase following liquid-gas phase transition. In 1990 Glendenning has pointed out that the phase equilibrium in multi-component system or in the case of plural chemical potential corresponding to the conserved quantities such as in neutron star matter must be carefully treated based on the more general Gibbs conditions \\cite{glen, gle92}. He emphasized that the phase transitions should be treated in a proper way in neutron-star matter, which is specified by two conserved quantities, electromagnetic charge as well as baryon number. In particular the local charge-neutrality condition used in the case of the Maxwell construction should be replaced by the global one in the mixed phase. Consequently the pressure becomes no more constant as in the Maxwell construction, but still increases in the mixed phase. He also demonstrated that the mixed phase is extended in the rather wide density region within a bulk calculation, where two sets of semi-infinite matter are considered in phase equilibrium, separated by a sharp boundary, and the pressure balance, chemical equilibrium and the global charge-neutrality conditions are imposed in the absence of the Coulomb interaction. It is important to see that particle fractions in each phase are no more the same to each other for given particle fractions as a whole. Nowadays we know that such phase transitions are rather common in other fields as well \\cite{mag96,ron,khra06,for07,net08,ios} and we call them {\\it non-congruent} phase transitions, following Iosilevskiy \\cite{ios}. Generally any phase transition in a system with two or more conserved charges must be non-congruent.  In this calculation the Coulomb interaction is discarded and we can by no means estimate how important it is, or the surface energy can not be taken into account. If we consider  a more realistic situation, we must consider inhomogeneous matter with various geometric structures called pasta, and take into account the finite-size effects such as the Coulomb interaction and the surface tension. Heiselberg et al.\\ have demonstrated that such finite-size effects are important for the mixed phase in the context of hadron-quark transition \\cite{HPS93}. Treating the pasta structure as in ref.~\\cite{Rav83}, they simply considered uniform particle densities in both phases. They showed a possibility of that the finite-size effects may largely restrict the pasta region. Subsequently, Voskresensky et al.\\ have shown that the rearrangement of particle densities and the screening effect for the Coulomb interaction becomes important, and  the large surface tension gives the mechanical instability of the geometrical structures in the mixed phase \\cite{voskre}. Following this idea, we have studied the pasta structures brought by the kaon condensation and hadron-quark deconfinement transition \\cite{marukaon06,end05,hyp07}. More recently we have extended our framework to deal with the finite temperature case or the neutrino-trapping matter, which may be relevant to supernova explosions or neutron star mergers \\cite{yas09}.  In this chapter we first present thermodynamic concepts of the non-congruent transition by considering the liquid-gas phase transition in asymmetric nuclear matter (section \\ref{cong}). In section \\ref{secLG} some results will be presented by a numerical calculation about the pasta structures in asymmetric nuclear matter. In section \\ref{secHQ} we will present the results about hadron-quark deconfinement transition, where the effects of finite temperature or neutrino trapping are also discussed. Section \\ref{sum} is devoted to summary and concluding remarks. ",
        "conclusions": "\\label{sum}  We have discussed general features of the phase transitions in compact stars, where multi-component substance are in equilibrium. These phase transitions are non-congruent, and the usual Maxwell construction cannot be applied any more. %show interesting phase diagrams. Using the bulk calculation, we have presented basic concepts of the non-congruent transition. Although these concepts are helpful in discussing and analyzing the properties of the phase transitions in compact stars, we must bear in mind that they are obtained in an idealized situation and must be carefully applied in realistic calculations by taking into account rather complicated effects such as the finite-size effects. Actually we found that uniform nuclear matter can exist in the mechanically unstable region.  When the finite-size effects are taken into account, the structured mixed phase (pasta phase) emerges to form the crystalline structures. Charge screening effect and the rearrangement effect of particle densities in the presence of the Coulomb interaction is also emphasized; they sometimes cause an instability of the geometrical structures in the mixed phase. In the extreme case the EOS resembles the one given by the Maxwell construction.  Liquid-gas transition in asymmetric nuclear matter and hadron-quark deconfinement transition are studied and their mixed phases have been calculated by using the Thomas-Fermi approximation for particle densities and the Wigner-Seitz (WS) cell approximation for the pasta structure. We have shown many interesting aspects of these phase transitions. However, we have also seen that the maximum mass of hybrid stars is rather low in the light of a recent observation \\cite{dem}, since contamination of hyperons considerably softens EOS. The transition to quark matter suppresses their appearance, but the maximum mass is still low due to the rather soft EOS of quark matter. So, we need another idea about EOS's of hadrons and quarks to circumvent this situation.  So far, most works on the pasta structures have used the WS cell with ansatz about the geometrical structures like droplet, rod, slab and so on. If we want to see the possibility of the intermediate shapes, or study the mixed phase in a more realistic way, a study is desirable without the WS cell approximation or ansatz about the geometrical structures. Recently we have performed a three-dimensional calculation of non-uniform nuclear matter based on the relativistic mean-field model and the Thomas-Fermi approximation \\cite{oka}. Introducing a cubic cell with periodic boundary conditions and discretizing it into grid, we numerically solve the coupled field equations for meson mean-fields and the Coulomb potential. Randomly distributing fermions (n, p, e) on the grid points as an initial condition, we relax their density distributions to attain the uniformity of the chemical potentials. As a preliminary result we demonstrate that pasta structures, previously given by the WS cell approximation, also appear in the three-dimensional calculation.   Formation of nuclear pasta has been discussed in supernovae \\cite{wata09} or possible observation of the signals of the mixed phase has been suggested in the spectra of the gravitational waves \\cite{sot}. The phenomenological implications of the pasta structures should deserve more elaborate studies; their thermodynamical or dynamical property is interesting; e.g. the neutrino opacity or the thermal conductivity may be affected by the pasta structures, and their elasticity may cause a discontinuous change of internal structure of compact stars. In particular the crystalline structure in the core region may bring about novel dynamical effects on compact-star phenomena by way of deformation or oscillation."
    },
    "1107/1107.3835_arXiv.txt": {
        "abstract": "Life arose on Earth sometime in the first few hundred million years after the young planet had cooled to the point that it could support water-based organisms on its surface.  The early emergence of life on Earth has been taken as evidence that the probability of abiogenesis is high, if starting from young-Earth-like conditions.  We revisit this argument quantitatively in a Bayesian statistical framework.  By constructing a simple model of the probability of abiogenesis, we calculate a Bayesian estimate of its posterior probability, given the data that life emerged fairly early in Earth's history and that, billions of years later, curious creatures noted this fact and considered its implications.  We find that, given only this very limited empirical information, the choice of Bayesian prior for the abiogenesis probability parameter has a dominant influence on the computed posterior probability.  Although terrestrial life's early emergence provides evidence that life might be common in the Universe if early-Earth-like conditions are, the evidence is inconclusive and indeed is consistent with an arbitrarily low intrinsic probability of abiogenesis for plausible uninformative priors.  Finding a single case of life arising independently of our lineage (on Earth, elsewhere in the Solar System, or on an extrasolar planet) would provide much stronger evidence that abiogenesis is not extremely rare in the Universe. ",
        "introduction": "\\label{sec:intro} Astrobiology is fundamentally concerned with whether extraterrestrial life exists and, if so, how abundant it is in the Universe.  The most direct and promising approach to answering these questions is surely empirical, the search for life on other bodies in the Solar System \\cite{chyba+hand2005, desmarais+walter1999} and beyond in other planetary systems \\cite{desmarais_et_al2002, seager_et_al2005}. Nevertheless, a theoretical approach is possible in principle and could provide a useful complement to the more direct lines of investigation.  In particular, if we knew the probability per unit time and per unit volume of abiogenesis in a pre-biotic environment as a function of its physical and chemical conditions and if we could determine or estimate the prevalence of such environments in the Universe, we could make a statistical estimate of the abundance of extraterrestrial life.  This relatively straightforward approach is, of course, thwarted by our great ignorance regarding both inputs to the argument at present.  There does, however, appear to be one possible way of finessing our lack of detailed knowledge concerning both the process of abiogenesis and the occurrence of suitable pre-biotic environments (whatever they might be) in the Universe.  Namely, we can try to use our knowledge that life arose at least once in an environment (whatever it was) on the early Earth to try to infer something about the probability per unit time of abiogenesis on an Earth-like planet without the need (or ability) to say how Earth-like it need be or in what ways.  We will hereinafter refer to this probability per unit time, which can also be considered a rate, as $\\lambda$ or simply the ``probability of abiogenesis.''  Any inferences about the probability of life arising (given the conditions present on the early Earth) must be informed by how long it took for the first living creatures to evolve.  By definition, improbable events generally happen infrequently.  It follows that the duration between events provides a metric (however imperfect) of the probability or rate of the events.  The time-span between when Earth achieved pre-biotic conditions suitable for abiogenesis plus generally habitable climatic conditions \\cite{kasting_et_al1993, selsis_et_al2007, spiegel_et_al2008} and when life first arose, therefore, seems to serve as a basis for estimating $\\lambda$. Revisiting and quantifying this analysis is the subject of this paper.  We note several previous quantitative attempts to address this issue in the literature, of which one \\cite{carter1983} found, as we do, that early abiogenesis is consistent with life being rare, and the other \\cite{lineweaver+davis2002} found that Earth's early abiogenesis points strongly to life being common on Earth-like planets (we compare our approach to the problem to that of \\cite{lineweaver+davis2002} below, including our significantly different results).\\footnote{There are two unpublished works (\\cite{brewer2008} and \\cite{korpela2011}), of which we became aware after submission of this paper, that also conclude that early life on Earth does not rule out the possibility that abiogenesis is improbable.} Furthermore, an argument of this general sort has been widely used in a qualitative and even intuitive way to conclude that $\\lambda$ is unlikely to be extremely small because it would then be surprising for abiogenesis to have occurred as quickly as it did on Earth \\cite{vanzuilen_et_al2002, westall2005, moorbath2005, nutman+friend2006, buick2007, sullivan+baross2007, sugitani_et_al2010}.  Indeed, the early emergence of life on Earth is often taken as significant supporting evidence for ``optimism\" about the existence of extra-terrestrial life ({\\it i.e.}, for the view that it is fairly common) \\cite{ward+brownlee2000, darling2001, lineweaver+davis2002}.  The major motivation of this paper is to determine the quantitative validity of this inference.  We emphasize that our goal is {\\it not} to derive an optimum estimate of $\\lambda$ based on all of the many lines of available evidence, but simply to evaluate the implication of life's early emergence on Earth for the value of $\\lambda$. ",
        "conclusions": "\\label{sec:conc} Within a few hundred million years, and perhaps far more quickly, of the time that Earth became a hospitable location for life, it transitioned from being merely habitable to being inhabited.  Recent rapid progress in exoplanet science suggests that habitable worlds might be extremely common in our galaxy \\cite{vogt_et_al2010, borucki_et_al2011, howard_et_al2011, wordsworth_et_al2011}, which invites the question of how often life arises, given habitable conditions.  Although this question ultimately must be answered empirically, via searches for biomarkers \\cite{kaltenegger+selsis2007} or for signs of extraterrestrial technology \\cite{tarter2001}, the early emergence of life on Earth gives us some information about the probability that abiogenesis will result from early-Earth-like conditions.\\footnote{We note that the comparatively very {\\it late} emergence of radio technology on Earth could, analogously, be taken as an indication (albeit a weak one because of our single datum) that radio technology might be rare in our galaxy.}  A Bayesian approach to estimating the probability of abiogenesis clarifies the relative influence of data and of our prior beliefs. Although a ``best guess'' of the probability of abiogenesis suggests that life should be common in the Galaxy if early-Earth-like conditions are, still, the data are consistent (under plausible priors) with life being extremely rare, as shown in Figure~3.  Thus, a Bayesian enthusiast of extraterrestrial life should be significantly encouraged by the rapid appearance of life on the early Earth but cannot be highly confident on that basis.  Our conclusion that the early emergence of life on Earth is consistent with life being very rare in the Universe for plausible priors is robust against two of the more fundamental simplifications in our formal analysis.  First, we have assumed that there is a single value of $\\lambda$ that applies to all Earth-like planets (without specifying exactly what we mean by ``Earth-like'').  If $\\lambda$ actually varies from planet to planet, as seems far more plausible, anthropic-like considerations imply planets with particularly large $\\lambda$ values will have a greater chance of producing (intelligent) life and of life appearing relatively rapidly, {\\it i.e.}, of the circumstances in which we find ourselves.  Thus, the information we derive about $\\lambda$ from the existence and early appearance of life on Earth will tend to be biased towards large values and may not be representative of the value of $\\lambda$ for, say, an ``average'' terrestrial planet orbiting within the habitable zone of a main sequence star.  Second, our formulation of the problem analyzed in this paper implicitly assumes that there is no increase in the probability of intelligent life appearing once $\\delta t_{\\rm evolve}$ has elapsed following the abiogenesis event on a planet.  A more reasonable model in which this probability continues to increase as additional time passes would have the same qualitative effect on the calculation as increasing $\\delta t_{\\rm evolve}$.  In other words, it would make the resulting posterior distribution of $\\lambda$ even less sensitive to the data and more highly dependent on the prior because it would make our presence on Earth a selection bias favoring planets on which abiogenesis occurred quickly.  We had to find ourselves on a planet that has life on it, but we did not have to find ourselves (\\emph{i}) in a galaxy that has life on a planet besides Earth nor (\\emph{ii}) on a planet on which life arose multiple, independent times.  Learning that either (\\emph{i}) or (\\emph{ii}) describes our world would constitute data that are not subject to the selection effect described above.  In short, if we should find evidence of life that arose wholly idependently of us -- either via astronomical searches that reveal life on another planet or via geological and biological studies that find evidence of life on Earth with a different origin from us -- we would have considerably stronger grounds to conclude that life is probably common in our galaxy.  With this in mind, research in the fields of astrobiology and origin of life studies might, in the near future, help us to significantly refine our estimate of the probability (per unit time, per Earth-like planet) of abiogenesis."
    },
    "1107/1107.4554_arXiv.txt": {
        "abstract": "We present a 3D photoionization model of the planetary nebula NGC~6302, one of the most complex and enigmatic objects of its kind. It's highly bipolar geometry and dense massive disk, coupled with the very wide range of ions present, from neutral species up to Si$^{8+}$, makes it one of the ultimate challenges to nebular photoionization modelling.  Our {\\sc mocassin} model is composed of an extremely dense geometrically thin circumstellar disk and a large pair of diffuse bipolar lobes, a combination which was necessary to reproduce the observed emission-line spectrum. The masses of these components, 2.2~M$_\\odot$ and 2.5~M$_\\odot$ respectively, gives a total nebular mass of 4.7~M$_\\odot$, of which 1.8~M$_\\odot$ (39\\%) is ionized. Discrepancies between our model fit and the observations are attributed to complex density inhomogeneities in the nebula. The potential to resolve such discrepancies with more complex models is confirmed by exploring a range of models introducing small-scale structures. Compared to solar abundances helium is enhanced by 50\\%, carbon is slightly subsolar, oxygen is solar, and nitrogen is enhanced by a factor of 6. These all imply a significant 3rd dredge-up coupled with hot-bottom burning CN-cycle conversion of dredged-up carbon to nitrogen. Aluminium is also depleted by a factor of 100, consistent with depletion by dust grains.  The central star of NGC~6302 is partly obscured by the opaque circumstellar disk, which is seen almost edge-on, and as such its properties are not well constrained. However, emission from a number of high-ionization `coronal' lines provides a strong constraint on the form of the high-energy ionizing flux. We model emission from the central star using a series of stellar model atmospheres, the properties of which are constrained from fits to the high-ionization nebular emission lines. Using a solar abundance stellar atmosphere we are unable to fit all of the observed line fluxes, but a substantially better fit was obtained using a 220,000~K hydrogen-deficient stellar atmosphere with $\\log g = 7.0$ and $L_\\star = 14,300$~L$_\\odot$. The H-deficient nature of the central star atmosphere suggests that it has undergone some sort of late thermal pulse, and fits to evolutionary tracks imply a central star mass of 0.73--0.82~M$_\\odot$. Timescales for these evolutionary tracks suggest the object left the top of the asymptotic giant branch $\\sim$2100 years ago, in good agreement with studies of the recent mass-loss event that formed one pair of the bipolar lobes. Based on the modelled nebular mass and central star mass we estimate the initial mass of the central star to be 5.5~M$_\\odot$, in approximate agreement with that derived from evolutionary tracks. ",
        "introduction": "Planetary nebulae (PNe) are one of the last evolutionary stages of the majority of low and intermediate mass stars (1-8~M$_{\\odot}$). During the previous asymptotic giant branch (AGB) evolutionary phase the star ejects a large fraction of its mass during brief phases ($10^3 - 10^4$~yrs) of high mass-loss (up to $10^{-4}$~M$_{\\odot}$~yr$^{-1}$). This mass-loss eventually exposes the central core of the star that radiates strongly in the UV and ionizes the surrounding nebula making it visible for study. Despite supposed spherically symmetric mass-loss while on the AGB, many PNe show point- or axi-symmetric shapes and complex morphologies that have been interpreted as evidence for the influence of a binary companion \\citep[e.g.][]{dema09}. This and other outstanding issues in PNe research such as their small-scale structures \\citep[e.g.][]{gonc01,mats09} and discrepancies between forbidden line and recombination line abundances for heavy elements \\citep{liu00,tsam04,wess05} suggest that out understanding of the late stages of stellar evolution is far from complete. Despite this, PNe are still regularly used as standard candles, metallicity indicators and tracers of stellar populations in distant galaxies \\cite[e.g.][]{buzz06}. If PNe are to be used in this way our understanding of their structure and evolution must be greatly improved, both through studies of large samples of such objects \\citep[e.g.][]{viir09,sabi10}, their precursors the AGB stars \\citep[e.g.][]{vand88,zijl01,wrig08}, and through individual case studies to understand the most complex examples \\citep[e.g.][]{gonc06,ware06}.  \\begin{figure} \\begin{center} \\includegraphics[width=180pt]{ngc6302.eps} \\caption{The planetary nebula NGC~6302 as imaged with the Wide Field Camera on the Hubble Space Telescope (HST). Blue is He~{\\sc ii}, brown is H$\\alpha$, and white is [S~{\\sc ii}]. The image is approximately $2 \\times 2.5$~arcmin or $0.6 \\times 0.75$~pc at the distance of NGC~6302. Image copyright: NASA and ESA.} \\label{ngc6302} \\end{center} \\end{figure}  NGC~6302 is one of the most complex and extreme examples of a planetary nebula. The nebula is broadly bipolar, with a second narrower pair of lobes visible further out (see Figure~\\ref{ngc6302}). It has a highly pinched waist that is typical of `butterfly' bipolar PNe \\citep{bali02}, as is a highly complex structure with many clumps and knots. The bipolar structure has been attributed to the presence of a very dense circumstellar disk that is seen almost edge-on and almost completely obscures the central star \\citep[$A_V \\sim 6.6$,][]{mats05,szys09}, which has only recently been detected. The disk radiates strongly in the infrared \\citep[e.g.][]{lest84} with evidence for significant fractions of very large grains \\citep{hoar92} and crystalline silicates \\citep{mols01}. NGC~6302 has been classified as a Type~I PN, both according to the criteria or \\citet[][N/O $\\geq$ 0.5 {\\em and} He/H $\\geq$ 0.125]{peim83} and the criteria of \\citet{king94} for a Milky-Way Type~I PN, that N/O $\\geq$ 0.8. The high helium and nitrogen abundances observed in the nebula imply a potentially massive progenitor, while observations of emission lines from very high ionization stages \\citep[e.g.][]{ashl88,rowl94,casa00} suggest an extremely hot central star.  Attempts to understand NGC~6302 have been hindered by its complexity and the inability to study its central star and therefore determine its evolutionary history. \\citet{alle81} could find ``no uniform density theoretical model that could reproduce the line intensities in NGC~6302 in any satisfactory fashion\", while \\citet{lame91} could not produce a physically plausible model for the nebula without requiring shock-excitation, which \\citet{orig93} and \\citet{casa00} found strong evidence against. However, many of these models were forced to assume spherical symmetry due to the computational constraints of full 3-dimensional modeling, yet the inhomogeneous morphology of NGC~6302 clearly requires a 3-dimensional model if reliable conclusions are to be drawn.  In this paper we present a 3D photoionization model of NGC~6302 using the 3D Monte Carlo photoionization and radiative transfer code {\\sc mocassin} \\citep{erco03,erco05,erco08}. The model presented here is a fully 3-dimensional model of NGC~6302 that includes both gas and dust, though we will only discuss the gas photoionization model in this paper. We will present the dust radiative transfer model in a future paper. In Section~2 we discuss the changes made to the {\\sc mocassin} code to enable the modeling of this object and outline the observations that our models are to reproduce. In Section~3 we describe the model used and the parameters used to build the model. In Section~4 we describe the modeling process and how different quantities varied according to changes in different parameters. In Section~5 we present the model results, including fits to the observed emission-line spectrum and the ionization and temperature structure of the nebula. Finally, in Section~6 we describe the resulting nebula properties and discuss the evolutionary history of the object. ",
        "conclusions": "We have constructed a 3D photoionization model of the extreme planetary nebula NGC~6302 using the Monte Carlo photoionization and radiative transfer code {\\sc mocassin}. Our model reproduces the majority of emission line strengths and ratios, which place strong constraints on many of the stellar and nebular properties. The model consists of a very dense circumstellar disk where densities reach 300,000~cm$^{-3}$ at the inner edge, a large pair of bipolar lobes with a constant density of 2000~cm$^{-3}$, and an intermediate component we have dubbed `the outflow'. This combination of components was required to match the majority of density- and temperature-sensitive line ratios from which a wide range of densities and temperatures have been observationally determined by previous authors. A number of line ratios remain not matched, which we suggest is due to complex density inhomogeneities throughout the nebula.  We derive a total nebular mass of 4.7~M$_\\odot$, of which 1.8~M$_\\odot$ is ionized, almost entirely in the bipolar lobes. The C/O ratio for NGC~6302 is 0.43 indicating a predominantly O-rich nebula, in agreement with all other measurements in the literature. Carbon is found to be marginally under-abundant compared to the solar value, while nitrogen is significantly over-abundant. The combined (C+N+O)/H abundance is 35\\% larger than the solar value, all of which suggests that a 3rd dredge-up occurred in the precursor of NGC~6302, enriching the central star in carbon and nitrogen, followed by hot bottom burning CN-cycle conversion of carbon to nitrogen.  In modeling the central star we have incorporated NLTE model atmospheres to reproduce the ionizing flux. Fits to the optical emission-line spectrum imply an extremely high temperature for the central star, and we used the observed emission from high ionization-stage infrared `coronal' lines to further constrain the form of the ionizing flux distribution. Using a solar abundance stellar atmosphere we were unable to fit all the observed line fluxes, overestimating many of the lower ionization stages while underestimating the higher stages. A substantially better fit was obtained using a hydrogen-deficient stellar atmosphere, implying that the central star of NGC~6302 is likely to be H-deficient.  Finally we compare the properties of the central star with evolutionary tracks for late thermal pulses that are capable of removing the majority of hydrogen from the central star. We find a good fit to a very late thermal pulse track for a star with an initial mass of 4--5~M$_{\\odot}$. Timescales for this evolutionary model imply that the central star left the top of the AGB $\\sim$2100 years ago, in good agreement with the age of one of the pairs of bipolar lobes determined by \\citet{meab08} to be $\\sim$2200 years. This mass is in reasonable agreement with the total modeled nebular mass plus central star mass of 5.5~M$_\\odot$, taking into account contributions from the dust component. The mass is also in agreement with the chemical abundances in the nebula such as the high helium abundance, low C/O ratio and slightly enhanced (C+N+O)/H ratio.  A future paper will discuss the dust component of this model that used {\\sc mocassin}'s ability to fully model the radiative transfer of dust in an ionized nebula. The model includes the majority of observed amorphous and crystalline dust species, using multiple dust chemistry distributions, non-standard grain size distributions, and a new light scattering code for non-spherical dust grains."
    },
    "1107/1107.3468_arXiv.txt": {
        "abstract": "It has been shown that anisotropy of homogeneous spacetime described by the general Kasner metric can be damped by quantum fluctuations coming from perturbative quantum gravity in one-loop approximation. Also, a formal argument, not limited to one-loop approximation, is put forward in favor of stability of isotropy in the exactly isotropic case. ",
        "introduction": "Standard \\citet{Jerusalem} and loop \\citet{ABL,Boj,AS} quantum cosmology heavily depends on the implicit assumption of (quantum) stability of general form of the metric. As a principal starting point in quantum cosmology, one usually chooses a metric of a particular (more or less symmetric) form. In the simplest, homogeneous and isotropic case, the metric chosen is the (flat) Friedmann--Lema\ufffdtre--Robertson--Walker (FLRW) one. Consequently, (field theory) quantum gravity reduces to a much more tractable quantum mechanical system with a finite number of degrees of freedom. It is obvious that such an approach greatly simplifies quantum analysis of cosmological evolution, but under no circumstances is it obvious to what extent is such an approach reliable. The quantum cosmology approach could be considered unreliable when (for example) the assumed symmetry of the metric would be unstable due to quantum fluctuations. More precisely, in the context of the stability, one can put forward the two, to some extent complementary, issues (questions): (1) assuming a small anisotropy in the almost isotropic cosmological model, have quantum fluctuations a tendency to increase the anisotropy or, just the opposite, to reduce it? (2) assuming we start quantum evolution from an exactly isotropic metric should be we sure that no quantum fluctuations are able to perturb the isotropy?  In this Letter, we are going to address the both issues of the quantum stability of spacetime metric in the framework of standard covariant quantum gravity. Namely, in Section 2, we address the first stability issue for an anisotropic (homogeneous) metric of the Kasner type, to one loop in perturbative expansion. In Section 3, we give a simple, formal argument, not limited to one loop, concerning the second issue. ",
        "conclusions": ""
    },
    "1107/1107.3142_arXiv.txt": {
        "abstract": "In a dwarf nova, the accretion disk around the white dwarf is a source of ultraviolet, optical, and infrared photons, but is never hot enough to emit X-rays.  Observed X-rays instead originate from the boundary layer between the disk and the white dwarf.  As the disk switches between quiescence and outburst states, the 2--10 keV X-ray flux is usually seen to be anti-correlated with the optical brightness.  Here we present \\xte\\ monitoring observations of two dwarf novae, VW Hyi and WW Cet, confirming the optical/X-ray anti-correlation in these two systems. However, we do not detect any episodes of increased hard X-ray flux on the rise (out of two possible chances for WW Cet) or the decline (two for WW Cet and one for VW Hyi) from outburst, attributes that are clearly established in SS Cyg.  The addition of these data to the existing literature establishes the fact that the behavior of SS Cyg is the exception, rather than the archetype as is often assumed.  We speculate on the origin of the diversity of behaviors exhibited by dwarf novae, focusing on the role played by the white dwarf mass. ",
        "introduction": "The outburst cycles of dwarf novae are understood as due to thermal instability in the accretion disk around the white dwarf primary. In the simplest form of the theory, the disk traces an S-curve in the surface density ($\\Sigma$)--effective temperature (\\teff) plane, with two stable branches.  The mass transfer rate from the Roche-lobe filling secondary is sufficiently low in dwarf novae that the disk is usually in the quiescent branch in which \\teff\\ is low and hydrogen is mostly neutral.  The disk is completely out of steady state, with accretion rate $\\dot m$ at each radius $r$ varying as $\\dot m(r)$ \\simpropto $r^3$, so the accretion rate onto the white dwarf is low and $\\Sigma$ increases over time.  When it reaches the end of the quiescent branch, the disk jumps to the outburst branch in which the disk is hot, ionized, and efficiently transports matter onto the white dwarf.  See \\cite{L2001} for a comprehensive review of the disk instability theory.  Dwarf novae are well known sources of 2--10 keV X-rays, particularly during quiescence but also in outburst.  Since the accretion disk surrounding a white dwarf is never hot enough to generate X-rays, even during outburst, we must seek their origin elsewhere. A widely held view is that they originate in an equatorial boundary \\citep{PR1985a,PR1985b}, which connects the Keplerian accretion disk to the slowly rotating white dwarf surface.  The detection of a narrow X-ray eclipse in OY Car (e.g., \\citealt{WW2003}) and other deeply eclipsing dwarf novae in quiescence is consistent with this picture, while casting doubt on the ``coronal siphon flow'' model of \\cite{MM1994}, which posits a hot, spatially extended X-ray emitting region around the white dwarf.  In the rest of this paper, we therefore proceed with the boundary layer model as the basis of our interpretation.  In this view, the quiescent boundary layer is optically thin and hot ($kT \\sim$10 keV) while the outburst boundary layer is predominantly optically thick with ($kT\\sim$10 eV) but also contains an optically thin region.  Coordinated multi-wavelength observations of dwarf novae can therefore reveal the physical states of the accretion disk and the boundary layer, at different points during the outburst cycle. Previous campaigns during a single outburst include those on SS Cyg \\citep{WMM2003}, VW Hyi \\citep{Wea1996,Hea1999}, and GW Lib \\citep{Bea2009}. These three campaigns in particular had dense enough X-ray coverage through the outburst to enable the authors to study the time evolution of X-ray flux.  Campaigns that covered multiple outburst cycles include those on YZ Cnc \\citep{VWM1999}, V1159 Ori \\citep{Sea1999}, and SU UMa \\citep{CW2010}. Here we present our analysis of two sets of X-ray observations taken with the PCA instrument \\citep{PCA} on board \\xte\\ \\citep{RXTE}, one of a superoutburst of VW Hyi, the other of WW Cet obtained over three months and containing two outbursts.  These data sets are well suited to the study of transitions of the X-ray emissions from quiescence to outburst and back. ",
        "conclusions": "This paper adds to the literature of X-ray observation of dwarf novae during outbursts.  Through these observations and analyses, common, if not universal, patterns are emerging.  In the majority of dwarf novae, X-ray flux is suppressed during outburst. This is true of 4 out of 6 systems listed in Table\\,\\ref{tab:paradigm}, probably also true in YZ Cnc and (largely) in V1159 Ori, as well as other dwarf novae for which one or two X-ray spectra have been obtained in outburst (e.g., Z Cam; see \\citealt{BWO2001}).  We interpret this as the consequence of an optically thick boundary layer, which radiates most of the boundary layer luminosity during outburst. In these majority, the X-ray spectrum also appears to become softer in outburst.  We propose Compton cooling as a possible mechanism for this, as \\cite{Nea2011} postulated for the quiescent X-ray emission from RS Oph.  Standard textbooks such as \\cite{FKR2002} convincingly show that Compton cooling cannot be important in cataclysmic variables if the seed photons are from the heated surface of the white dwarf in a 1-zone accretion. However, the boundary layer in dwarf novae in outburst has at least two zones. At the equator, the boundary layer is optically thick, emitting copious soft photons; at somewhat higher latitude, the boundary layer is optically thin but subject to these external seed photons (see Figure 8 of \\citealt{PR1985a}).  Since the number of soft photons greatly exceeds that of the hot electrons (e.g., by a factor of 1,000 if the energy of a 10 keV electron is radiated away by 10 eV photons), each electron can experience many Compton scattering events even though each photon scatters once at most, thus keeping the soft spectrum largely unchanged (i.e., no hard inverse Compton component).  Since the density of a cooling flow plasma rapidly increases towards lower temperatures, the optically thin region of the boundary layer is predominantly Compton cooled at high temperatures and predominantly Bremsstrahlung cooled at low temperatures.  The observed soft-to-hard X-ray luminosity ratio for SS Cyg in outburst is $\\sim 40$ \\citep{WMM2003}, which makes it plausible that the optical thin part of the boundary layer is partly Compton-cooled, since the two emitting regions lie close to each other within the boundary layer.  To interpret the X-ray transitions, we must consider the disk instability model in more detail.  The instability is principally a local property of the disk.  However, once a transition from the quiescent state to the outburst state happens in one part of the disk, this propagates in the form of a heating front, behind which the mass inflow rate becomes much higher. X-ray behaviors are altered when the heating front reaches the inner edge of the disk.   The optical flux, on the other hand, predominantly arises from the outer parts of the disk, which have the most emitting area.  It is likely that the outburst can start at a variety of radial distances from the white dwarf in different systems or in different outbursts, and this is described as ``outside-in'' and ``inside-out'' outbursts in the studies of UV delay (see, e.g., \\citealt{CWP1986}).  It is therefore possible to have a range of X-ray delay times among systems, and among different outbursts for a single dwarf nova. On the return to quiescence, the cooling front is likely to start from near the outer edge of the disk, and so X-ray transition happens when it reaches the boundary layer, near the end of optical decline.  From our own analysis and a survey of the literature, we conclude that the temporary enhancement of X-ray flux during the rise and late decline has only been definitively seen in SS Cyg so far.  There have only been two studies that attempt to incorporate the observational constraints imposed by the X-ray observations of dwarf novae throughout the course of an outburst  into the disk instability model: The calculations of \\cite{SHL2003} were specifically for SS Cyg, while those of \\cite{SHL2004} were tailored for VW Hyi. These papers adopt a critical value of ${\\dot M}_{\\rm crit} = 10^{16}$ g s$^{-1}$ in the inner disk as the dividing line between having an optically thin or thick boundary layer, following the empirical determination by \\cite{PR1985b}.  A comparison of the theoretical X-ray fluxes shown in Fig. 4 of \\cite{SHL2003} with those of Fig. 9 of \\cite{SHL2004} shows the X-ray enhancements, both during rise and decline, to be common to both models. Thus, the model can account for the behavior exhibited by SS Cyg (see Fig. 2 of \\citealt{WMM2003}), but not those of SU UMa, VW Hyi, or WW Cet (Table 1).  To investigate this further, we have derived the critical accretion rate observationally for the 6 systems listed in Table\\,\\ref{tab:paradigm}. We assume that the highest luminosity observed for the optically thin component reflects (or is close to) this critical value.  We have collected the information in Table\\,\\ref{tab:mcrit}, including the adopted values and references for the distances and the white dwarf mass.  We convert the bolometric X-ray luminosity to accretion rate using $L=GM_1 \\dot m/2R$, for a range of plausible values of M$_1$.  We also plot the results in Figure\\,\\ref{mcrit} against M$_1$.  For five systems, the plots are limited to the range of M$_1$ inferred from the reference cited. For SU UMa, we show three parallel lines over the entire mass range of the figure, as there are no published estimates for M$_1$. The objects are color-coded based on the presence or otherwise of hard X-ray suppression and temporary rise.  For comparison, we also plot the theoretical curves from Figure\\,8(a) of \\cite{PN1995}.  Unfortunately, we are unable to reach a firm conclusion, except that the observationally derived values are all much lower than what \\cite{PN1995} predicted.  This may not be a huge problem, since the specific values depend on the assumptions made in the study, including the adopted value for the viscosity parameter, $\\alpha$.  The authors themselves caution against adopting these precise theoretical values, although the strong dependence on M$_1$ is likely to be a more robust conclusion.  The values we derived observationally are in closer agreement with the empirical conclusion of \\cite{PR1985b}.  Nevertheless, we propose that the diversity of behavior shown by these 6 dwarf novae reflects the fact that ${\\dot M}_{\\rm crit}$ is in fact a function of various system parameters, and not uniform among all dwarf novae. In this view, the quiescent accretion rate is already close to this value in some systems, in which case any increase during outburst will immediately lead to the development of an optically thick boundary layer and hard X-ray suppression.  Moreover, we propose that the primary mass is likely to be a key parameter, partly because the conversion from luminosity to accretion rate depends on $M_1$, and partly because the theoretical investigation by \\cite{PN1995} indicates that the critical rate depends steeply on M$_1$.  Looking at individual systems, the inferred critical accretion rate for VW Hyi is extremely low, compared with those for other systems as well as theoretical expectations.  We note, however, that the distance estimate for this system relies on an indirect argument without any indication of the size of the error bar, and should be refined.    As for the two systems that do not show X-ray suppression during outburst, GW Lib can be understood as due to the extremely low accretion rate in quiescence \\citep{Hea2007}.  Interestingly, the inferred accretion rate for U Gem is also low.  If the current mass estimate is reliable, it also has the most massive white dwarf of the 6 systems, which suggests that the critical accretion rate is also highest.  However, both these systems are reported to have developed an optically thick boundary layer, so the lack of hard X-ray suppression is a puzzle, particularly in light of our Compton-cooling interpretation for systems that do show hard X-ray suppression.  We believe that we should no longer assume a single critical accretion rate for all dwarf novae.  Rather, we propose that future efforts be focused on determining which system parameters determine the critical accretion rate.  Specifically, X-ray observations of more dwarf novae with well-determined distances and white dwarf masses should be obtained through outburst cycles.  In particular, systems with extreme masses should be high priority targets, given the likely steep dependence of the critical accretion rate on the white dwarf mass."
    },
    "1107/1107.0141_arXiv.txt": {
        "abstract": "We investigate the properties of satellite galaxies that surround isolated hosts within the redshift range $0.01 < z < 0.15$, using data taken as part of the Galaxy And Mass Assembly survey. Making use of isolation and satellite criteria that take into account stellar mass estimates, we find 3\\,514 isolated galaxies of which 1\\,426 host a total of 2\\,998 satellites. Separating the red and blue populations of satellites and hosts, using colour-mass diagrams, we investigate the radial distribution of satellite galaxies and determine how the red fraction of satellites varies as a function of satellite mass, host mass and the projected distance from their host. Comparing the red fraction of satellites to a control sample of small neighbours at greater projected radii, we show that the increase in red fraction is primarily a function of host mass. The satellite red fraction is about 0.2 higher than the control sample for hosts with $11.0 < \\log_{10} {\\mathcal M}_{*} < 11.5$, while the red fractions show no difference for hosts with $10.0 < \\log_{10} {\\mathcal M}_{*} < 10.5$. For the satellites of more massive hosts the red fraction also increases as a function of decreasing projected distance. Our results suggest that the likely main mechanism for the quenching of star formation in satellites hosted by isolated galaxies is strangulation. ",
        "introduction": "\\label{intro}  In recent years satellite galaxies have received much attention in both observational and theoretical studies in order to establish their role in the formation and evolution of galaxies. In the current $\\Lambda$CDM models of the Universe, galaxies are assembled in a hierarchical fashion, whereby small halos of dark matter (DM) merge to form larger halos, in which baryonic matter then cools and condenses to form stars. In this framework satellite galaxies are associated with sub-halos of dark matter residing within the virial radii of larger halos, which are believed to be left over from an earlier assembly phase of their host. Thus the measurement of the spatial distribution of satellites, both in terms of their angular and radial distributions, can provide an insight into the mass accretion histories of galaxies.  The most common way to determine the radial distribution is to calculate the projected density of satellites surrounding samples of more luminous hosts, which requires redshifts and photometry of satellites and hosts. Previous studies which attempted to constrain the small-scale galaxy correlation function \\citep{Lake1980,Phillips1987,Vader1991,Lorrimer1994} or investigated the companions of field ellipticals \\citep{Madore2004, Smith2004} found that the projected density of satellites, as a function of radius, has a profile that can be described as a power law of the form $\\Sigma(R) \\propto R^{\\alpha}$, with a slope, $\\alpha$ ranging between $-0.5$ and $-1.25$. Subsequently, more accurate measurements of the projected density have been made, making use of larger and more complete redshift surveys, such as the 2-degree Field Galaxy Redshift Survey (2dFGRS, \\citealt{Colless2001}) and the Sloan Digital Sky Survey (SDSS, \\citealt{York2000}), allowing the production of large statistical samples of satellites and hosts selected using various well defined criteria \\citep{Sales2004,Chen2006, Ann2008,Chen2008, Bailin2008}.  Studies which have examined the radial distribution of satellites as a function of host luminosity, colour or morphology, have obtained mixed results. The first to divide their sample into early and late types was \\cite{Lorrimer1994}, who found the projected density profile of the satellites of early-type hosts to have steeper slopes and therefore to be more centrally concentrated than the satellites of late-type hosts. In contrast to this, \\cite{Sales2005} using data from 2dFGRS, found that the distribution of the satellites of red hosts has a shallower profile than the satellites of blue hosts, which even deviates from a power law, flattening at small projected separation. More recently \\cite{Chen2008} using SDSS data, found that the satellites of both red and blue hosts follow similar power slopes after correcting for interlopers (galaxies mistaken as satellites through projection but not actually physically bound to their host).  Dividing their satellite sample into red and blue populations, \\cite{Chen2008} also finds that red satellites are more centrally concentrated than blue satellites, a trend which is also seen in the semi-analytic galaxy samples produced by \\cite{Sales2007} using the Millennium Simulation. One explanation of this is that red satellites were accreted into their host's halo at earlier times than the blue satellites, which is consistent with the observational finding that red satellites have an anisotropic angular distribution with a preference of being aligned along the major axes of their hosts \\citep{Brainerd2005, Yang2006,Azzaro2007,Bailin2008,Agustsson2010}. Comparisons between the radial distribution of satellites and the dark matter distribution produced from simulations have also been conducted by \\cite{Chen2006} and \\cite{VDBosch2005}, who find that satellites are more centrally concentrated than dark matter sub-halos, but consistent with the dark matter profile.  Satellite galaxies not only provide useful information about the formation of galaxies but also about the processes that govern galaxy evolution in localised environments on scales of $\\sim 1$\\,Mpc. In the current theory of galaxy evolution it is thought that virtually all galaxies start off as blue, late-type discs which are then transformed by various processes into red, early types. This is supported by the observed bimodality of galaxies which can be seen out to $z \\sim 1$ \\citep{Bell2004, Willmer2006, Prescott2009}, and studies such as \\cite{Willmer2006} and \\cite{Faber2007} which have shown there has been a doubling in the stellar mass density of galaxies on the red sequence over the last 7-8 Gyr.  The main processes believed to be responsible for transforming blue/late-type galaxies into red/early types involve the quenching of star formation, and include major-mergers \\citep{Toomre1972, Hopkins2008a}, feedback from active galactic nuclei \\citep{Bower2006,Croton2006} and the depletion of gas reservoirs that fuel star formation. For satellite galaxies, the dominant process is most likely to be gas depletion, caused by the stripping of gas via a number of different hydrodynamical and radiative interactions with their hosts, acting over different timescales.  {\\revone When a satellite halo is accreted by the larger halo of its host, hot gas from the satellite may be removed in the process known as strangulation \\citep{Larson1980, Balogh2000}, resulting in the gradual decline in star formation over long timescales ($> 1$ Gyr), as its fuel for future star formation is depleted. Star formation can be shut off more rapidly if the satellite is subjected to sufficient external pressure that its cold gas reservoir is removed in the process of ram-pressure stripping \\citep{Gunn1972}. Gas stripping via harassment \\citep{Moore1996} whereby the dark matter sub-halos of satellites are heated after undergoing frequent high-velocity encounters with other dark matter halos, is also a possibility, although this process is more likely to occur in galaxy clusters rather than small groups.  Recent studies using the SDSS have indicated that strangulation is the main mechanism causing the transition of satellites from the blue to the red sequence. Producing a group catalogue from DR2, \\cite{Weinmann2006} investigate how the fractions of early- and late-type satellites vary as a function of halocentric radius, halo mass and luminosity, {\\revtwo observing that the early-type fraction increases with decreasing halocentric radius, increasing halo mass and increasing luminosity. They argue that the increase in early-type fraction with luminosity at fixed halo mass is not expected if ram-pressure stripping or harassment is the primary cause of gas removal.} \\cite{Ann2008} use SDSS DR5 to investigate how the early-type fractions of satellites surrounding isolated hosts vary as a function of luminosity and projected distance. They find that the early-type satellite fraction increases significantly with decreasing projected radius for early-type hosts and stays approximately constant for late types. They conclude that hot X-ray emitting gas of the early-type hosts is responsible for the {\\revtwo removal} of gas.  Using an SDSS DR4 group catalogue, \\cite{VDBosch2008} compare the concentrations and colours of centrals and satellites of the same mass. By matching central and satellite pairs in both stellar mass and concentration, they find there is a significant difference in colour. Centrals and satellites matched in both stellar mass and colour, on the other hand, show no difference in concentration. Under the assumption that centrals are the progenitors of satellites (centrals change into satellites after being accreted into a larger halo), this implies that either the strangulation or ram pressure stripping process is occurring rather than harassment, which is believed to have a significant effect on the morphology of galaxies. Investigating the red fraction of satellites by mass they estimate that 70 per cent of satellite galaxies with ${\\mathcal M}_{*} = 10^{9} {\\mathcal M}_{\\odot}$ have undergone satellite quenching in order to be on the red sequence at the present, with 30 per cent already red before becoming a satellite. For more massive satellites they find quenching to be less effective, with 65 per cent of satellites with ${\\mathcal M}_{*} = 10^{10} {\\mathcal M}_{\\odot}$ being red before accretion and virtually all satellites with ${\\mathcal M}_{*} = 10^{11} {\\mathcal M}_{\\odot}$ already being red before becoming satellites.}  Furthermore, satellites are believed to affect the evolution of their hosts. Minor mergers between dwarf satellites and their central hosts provide one way in which to distort and thicken galaxy discs \\citep{Quinn1993} and enlarge the bulge components of discs \\citep{Dominguez2008}, and mergers involving gas-rich satellites may also provide gas which could prolong or replenish star formation in early-type spirals \\citep{White1991,Hau2008}.  {\\revone Finally, satellite systems analogous to the Milky Way (MW)-Magellanic clouds systems have also become of much interest lately due to their apparent rarity. Observational studies such as \\cite{James2011}, \\cite{Liu2011} and \\cite{Tollerud2011} find that only $\\sim$10 per cent of MW like hosts have one Magellanic cloud like satellite and only $\\sim$5 per cent have two. Similarly N-body simulations have shown that less than 10 per cent of MW sized dark matter haloes contain two Magellanic cloud sized sub-halos \\citep{Busha2010,Boylan-Kolchin2010}.}  In this paper we use data taken from the Galaxy And Mass Assembly (GAMA) survey to investigate: firstly the radial distribution of satellites surrounding a sample of isolated host galaxies, as a function of host mass and colour; and secondly how the red fraction of satellites depends on the projected distance between satellite and host, the stellar mass of the satellite and the stellar mass of the host. {\\revtwo The GAMA redshift survey extends up to 2\\,mag deeper than the SDSS Main Galaxy Sample and, unlike previous studies, we use isolation criteria that take account of stellar mass estimates. The results are discussed in terms of} the quenching mechanisms that act on satellites.  The structure of this paper is as follows. In Section~2 we outline the GAMA survey from which our samples of galaxies are taken. In Section~3 we define our criteria used to select isolated hosts and satellite galaxies. We also define how we select a control sample of neighbours. Section~4 describes how we divide the red and blue populations of galaxies, compare properties of the hosts and satellites, and determine the projected density of satellites as a function of host mass and host colour. We show how the red fraction of satellites depends on satellite mass, host mass and projected radius and discuss potential processes that could produce these results in Section~5. Finally in Section~6 we summarise our main results. Throughout this paper we assume values of $H_0 = 70$ km\\,s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We have performed a search for the satellites of isolated galaxies from the GAMA survey using selection criteria that take into account stellar mass estimates. Separating galaxies into the red and blue populations using a colour-mass diagram, we have investigated the radial distribution of satellites, and how the red fraction of satellites varies as function of projected distance, host stellar mass and satellite stellar mass.  We find out of a sample of 3\\,514 isolated galaxies, 1\\,426 which host a total of 2\\,998 satellite galaxies. Separating the galaxies into red and populations, we find 41.2 per cent of the satellites are blue with red hosts, 33.3 per cent are red with red hosts, 22.2 per cent are blue with blue hosts and only 3.3 per cent are red with blue hosts. We find mean stellar masses and absolute magnitudes of $\\log_{10} {\\mathcal M}_{*} = 10.00$ and $M_{r} =-19.96$, and $\\log_{10} {\\mathcal M}_{*} = 9.35$ and $M_{r} =-19.30$, for the red and blue populations of satellites respectively. The average satellite in this sample is found to be $\\sim$ 1/10th of the mass of its host. The mean $\\Delta \\log_{10} {\\mathcal M}_{*} = 1.18$. Our main results are as follows:  \\begin{enumerate}  \\item Parametrizing the projected density of the satellites by $\\Sigma(R) \\propto R^{\\alpha}$, we find best fitting slopes of $\\alpha \\simeq -1.0$ for satellites divided into four host stellar mass bins (Fig.~\\ref{fig:rprojmass}). By dividing the satellite hosts into red and blue populations we find the satellites of blue, low mass hosts to be more centrally concentrated. Power law slopes of $\\alpha = -0.94$ and $\\alpha = -1.05$ for the satellites of red and blue hosts are obtained, respectively. Correcting for interlopers the slopes steepen to $\\alpha = -1.13$ and $\\alpha = -1.38$ (Fig.~\\ref{fig:rprojcol}).  \\item We find there to be a steady increase in the red fraction of satellites as projected distance decreases for satellites with hosts of stellar mass $\\log_{10} {\\mathcal M}_{*} \\ge 10.5$. The red fraction of satellites with hosts of stellar masses $\\log_{10} {\\mathcal M} < 10.5$ shows no trend as a function of radius (Fig.~\\ref{fig:rfhost}).  \\item Sub-dividing the satellite sample into host and satellite mass bins reveals that the host mass has the biggest effect on the satellite red fraction. Satellites with more massive hosts are more likely to be red. Comparing satellites to the small neighbours control sample, we find the red fraction is 0.13--0.25 higher for satellites of $11.0 \\le \\log_{10} {\\mathcal M}_{*} \\le 11.5$ hosts, 0.06--0.15 higher for $10.5 \\le \\log_{10} {\\mathcal M}_{*} \\le 11.0$ hosts, and almost unchanged for lower-mass $10.0 \\le \\log_{10} {\\mathcal M}_{*} \\le 10.5$ hosts (Fig.~\\ref{fig:rfmain}).  \\end{enumerate}  {\\revtwo The effect of environment, as noted by the difference between the satellites and small neighbours, can act over large range in high mass haloes and appears to be primarily a function of host mass and not satellite mass.}  These results suggest that quenching of star formation by strangulation is likely to be the main gas removal process in satellite systems, with other tidal effects, having a smaller contributing effect. Further study of the radial dependence of the instantaneous star formation rates of satellite galaxies using H$\\alpha$ emission line measurements, is likely to yield more details about the quenching mechanism."
    },
    "1107/1107.0970_arXiv.txt": {
        "abstract": "{A debris disk is a constituent of any planetary system surrounding a main sequence star. We study  whether close stellar encounters  can disrupt and  strip a debris disk of its planetesimals in the expanding open cluster of its birth with a decreasing star number density over 100~Myrs. Such stripping would affect the dust production and hence detectability of the disk. We tabulated the fractions of planetesimals stripped off  during stellar flybys of miss distances between 100 and 1000 AU and for several  mass ratios of the central to passing stars. We then  estimated  the numbers of close stellar encounters over the lifetime of several expanding open clusters characterized  by their initial star densities. We found that a standard disk, with inner and outer radii of 40 and 100~AU, suffers no loss of planetesimals over 100~Myrs around a star born in a common embedded cluster with star density $\\le 1000~pc^{-3}$. In contrast, we found that such a disk  is  severely depleted of its planetesimals over this timescale around a star born in an Orion-type  cluster where the  star density is $>20~000~pc^{-3}$. In this environment, a disk loses  $>97$\\% of its planetesimals around an M-dwarf, $>63$\\%  around a solar-type star, and $>42$\\% around an A-dwarf, over 100~Myrs. We roughly estimate that two-thirds of the stars may be born in such high star density clusters. This might explain in part why fewer debris disks are observed around lower mass stars. ",
        "introduction": "A debris disk surrounding a main sequence star is the collection of planetesimals (comets or asteroids) that are the leftovers from an early phase of planet formation. A debris disk is a constituent  of any planetary system in the core-accretion theory of  planet formation. It is the  analogue of  the Kuiper Belt   beyond Neptune, or of the asteroid belt between Mars and Jupiter, in our Solar System. The formation and evolution of debris disks and planetary systems are inter-related in theory, and can provide complementary insight into each other.    \\begin{figure*}[t] \\centering \\includegraphics[width=17.cm, angle=0, bb=0 300 595 842]{multiplt.ps} \\caption{Stellar flyby. The debris disk is non-self-gravitating, and has inner and outer radii of 40 and 100 AU at the start of the simulation. The central star is at the origin of the coordinate system. The trajectory of the passing star is parabolic, coplanar and  prograde with respect to the disk.  The mass ratio of the two stars is unity. At the closest approach, the miss distance  is 200 AU and the resulting maximum velocity is 4.2 km/s. The fraction of planetesimals stripped  off the disk and lost by the system is 13\\% during this close stellar encounter. \\label{flyby}} \\end{figure*}   The dust content of a debris disk is the key to its observability. Small dust grains scatter the light of the central star and  make the disk visible at optical wavelengths. Large grains are  heated by the central star and efficiently re-radiate in the far-IR and (sub)millimeter domains where they are observed. It is thought that dust must be continually or episodicly  replenished by  mutual collisions between planetesimals because dust grains are removed on short timescales (a few millions years or less). Disks have to be several-fold dustier than the Kuiper belt $(\\times 100$) to be detectable with our  instrumental sensitivities, even for the nearest stars. Nonetheless, bright  disks are detected  around $32\\pm5$\\% of A-type dwarfs (Su et al 2006, Wyatt et al 2003), $16\\pm2.8$\\% of solar-type dwarfs (Bryden et al 2006,  Trilling et al 2008, Najita \\& Williams 2005) and $\\le$5\\% of M-type dwarfs (Lestrade et al 2006 and 2009, Gautier et al 2007) according to searches for cold dust at  $\\lambda~=~70\\mu$m by Spitzer and at (sub)millimeter wavelengths by radiotelescopes. Debris disks were recently reviewed by  Wyatt (2008).  The mechanisms responsible for  enhanced collisional activity  and  for grinding planetesimals to dust  in debris disks remain unclear. The statistics of disk detections just recalled cannot be accurately interpreted  unless these mechanisms are identified. It has been proposed that unseen planets dynamically stir the disk and generate copious amount of dust that evolve under gravity, radiation pressure, P-R, and stellar-wind drags (Moro-Mart\\'in \\& Malhotra, 2002;  Wyatt  2003; Mustill \\& Wyatt 2009; Kennedy \\& Wyatt 2010). Steady-state  equilibrium between collisional cascades of planetesimals  and dust removal processes  can explain how   disks fade with age  (Dominik \\& Decin 2003). Orbit resonance crossing of giant planets can trigger abrupt showers of planetesimals in the system, severely depleting the disk and limiting its detectability (Gomes et al 2005, Morbidelli et al 2005, Tsiganis et al 2005).  However, the distribution of disk fractional luminosities with star ages is inconsistent with the idea that this mechanism is common among debris disks (Booth et al. 2010). Finally, distant icy-planets can successively form in waves outwards in the disk and generate dust rings (Kenyon \\& Bromley, 2002).  In this paper, we study whether stellar flybys  during the first 100 Myrs of the lifetime of a star, while it is still in the expanding open cluster of its birth, can significantly deplete a disk of its planetesimals, affecting its dust production and therefore  detectability.  Stellar flybys have been invoked to explain the high eccentricity orbits of some Kuiper Belt Objects such as  Sedna (Kenyon \\& Bromley, 2004), the dynamics of planetary systems (Malmberg et al. 2007, Spurzem et al. 2009), and the structures of  debris disks (Larwood, 1997; Mouillet et al. 1997;  Kalas et al 2000; Kobayashi \\& Ida,  2001). Impacts of stellar flybys during the first few million years have also been studied in the context of the evolution of protoplanetary disks and planet formation (Bonnell et al 2001,  Olczack, Pfalzner \\& Eckart 2009).  For our study, we did not resort  to a full N-body integration of an  expanding cluster with stars surrounded by disks of planetesimals. We instead divided the problem into parts by first estimating the fraction of planetesimals stripped off a disk by the passage of a star moving on a parabolic trajectory  in Sect.~2, and second by estimating the probability of close stellar encounters in an expanding cluster in Sect.~3. Finally, we discuss our results in Sect.~4. ",
        "conclusions": "We have studied the depletion of planetesimals in a debris disk  triggered by close stellar encounters in the environment of an expanding open cluster over its lifetime of 100~Myrs. We have found that depletion is significant only for an initial star-density  of the embedded cluster of origin greater than $20~000~pc^{-3}$, as in the Orion Nebula Cluster, and in a disk of standard size (inner and outer radii~: 40 and 100 AU). In these conditions, a debris disk loses $>$97\\% of its  planetesimals around an M-dwarf, $>$63\\%  around a solar-type star, and $>$42\\%  around an A-dwarf in 100 Myr. This level of depletion  could affect two-thirds of the stars searched in  surveys since two-thirds of them are born in  high star-density embedded clusters  according to  the catalog of Lada and Lada (2003). However, more compact disks (inner and outer radii~: 10 and 40~AU) are much less sensitive to their environment. Unfortunately, debris disk sizes are not yet well enough known  to decisively conclude whether the observed trend of fewer debris disks being detected around lower mass stars can be  explained by the mechanism we have studied."
    },
    "1107/1107.0836_arXiv.txt": {
        "abstract": "The classical method for measure the Earth-Sun distance is due to Aristarchus and it is based upon the measure of the angle Moon-Earth-Sun when the Moon is exactly in quadrature. Such an angle is only 9 arcminutes smaller than 90 degrees, and it is very difficult to evaluate, being necessary to look directly towards the Sun. The distance Earth-Moon and the Earth's diameter are necessary ingredients in order to derive the value of the astronomical unit. This method requires also the knowledge of the Moon's distance and the Earth's diameter, but it can permit a more precise measurement of the involved angles. ",
        "introduction": "\\begin{figure}% \\centerline{\\psfig{file=eclipse1.eps,width=12cm}} \\vspace*{0pt} \\caption{The geometry of the Earth shadow during a lunar eclipse.} \\end{figure}    During a total eclipse of the moon, and even during a partial eclipse of magnitude something larger than 1/2, it is possible to observe the configuration shown in figure 2, where the shadow of the Earth crosses the full Moon along a diameter. Taking a photo or projecting the image of the Moon with a pinhole\\cite{sig} on a white paper it is possible to draw with a pencil the Earth's shadow's profile, in order to evaluate its diameter with respect to that of the Moon. The same geometry was utilized by Hipparcos for evaluating the position of the centre of the Earth shadow during the eclipse, studying the  precession of the equinoxes.\\cite{bar}  From the distance DC expressed in term of the fraction $\\frac{1}{a}$ of the Moon radius $r_{Moon}=EC$, we obtain that the Earth's shadow radius at the Moon's distance $R_{\\oplus,M}$ is given by the equation: \\be R_{\\oplus,M} = r_{Moon} \\cdot \\frac{a^2+1}{2a} \\ee at the Moon's distance.  Such value is to compare with the Earth's diameter, already known by Eratostenes,\\cite{bar} in order to evaluate the angle with which the Sun is seen to appear as large as the Earth, as seen from the vertex of the shadow's cone. Also the position of that vertex is calculable.  \\begin{figure}% \\centerline{\\psfig{file=eclipse2.eps,width=12cm}} \\vspace*{0pt} \\caption{The geometry of the Earth's shadow above the eclipsed Moon.} \\end{figure}  The position of the Sun is somewhere in the backwards extension of shadow's cone. It needs another constraint in order to fix its position. That is possible knowing its angular extension, a measure obtainable with a daily pinhole projection with a precision up to 30 arcseconds.\\cite{sig} ",
        "conclusions": "That method requires the same hypothesis of the Aristarchus' one,\\cite{bar} and it requires instrumentations available even during classical epochs. So it can be considered equivalent to a classical method for measuring the Earth-Sun distance."
    },
    "1107/1107.2234_arXiv.txt": {
        "abstract": "Half a century has passed since the foundation of nuclear astrophysics. Since then, this discipline has reached its maturity. Today, nuclear astrophysics constitutes a multidisciplinary crucible of knowledge that combines the achievements in theoretical astrophysics, observational astronomy, cosmochemistry and nuclear physics. New tools and developments have revolutionized our understanding of the origin of the elements: supercomputers have provided astrophysicists with the required computational capabilities to study the evolution of stars in a multidimensional framework; the emergence of high-energy astrophysics with space-borne observatories has opened new windows to observe the Universe, from a novel panchromatic perspective; cosmochemists have isolated tiny pieces of stardust embedded in primitive meteorites, giving clues on the processes operating in stars as well as on the way matter condenses to form solids; and nuclear physicists have measured reactions near stellar energies, through the combined efforts using stable and radioactive ion beam facilities.  This review provides comprehensive insight into the nuclear history of the Universe and related topics: starting from the Big Bang, when the ashes from the primordial explosion were transformed to hydrogen, helium, and few trace elements, to the rich variety of nucleosynthesis mechanisms and sites in the Universe. Particular attention is paid to the hydrostatic processes governing the evolution of low-mass stars, red giants and asymptotic giant-branch stars, as well as to the explosive nucleosynthesis occurring in core-collapse and thermonuclear supernovae, $\\gamma$-ray bursts, classical novae, X-ray bursts, superbursts, and stellar mergers. ",
        "introduction": "Most of the ordinary (visible) matter in the Universe, from a tiny terrestrial pebble to a giant star, is composed of protons and neutrons (nucleons). These nucleons can assemble in a suite of different nuclear configurations. The masses of such bound nuclei are actually smaller than the sum of the masses of the free individual nucleons. The difference is a measure of the nuclear binding energy that holds the nucleus together, according to Einstein's famous equation $E=\\Delta mc^2$, where $c$ is the speed of light and $\\Delta m$ denotes the mass difference between the bound nuclei and the free nucleons. For example, the nuclide $^{56}$Fe is among the most tightly bound species, with a binding energy near 8.8 MeV per nucleon. There are 82 elements that have stable isotopes, all the way from H to Pb (except for Tc and Pm). Several dozen elements have only unstable isotopes, naturally abundant or artificially produced in nuclear physics laboratories (the super-heaviest nucleus discovered to date has 118 protons and a half-life of 0.89 ms \\cite{Oga07}). The distribution of the nuclide abundances in the Solar System, plotted as a function of the number of nucleons (mass number), $A$, is shown in Fig. \\ref{SSabun}. It displays a complicated pattern, with H and He being by far the most abundant species. The abundances generally decrease for increasing mass number. Exceptions to this behavior include the light elements Li, Be and B, which are extremely underabundant with respect to their neighbors, a pronounced peak near $^{56}$Fe (the so-called {\\it iron-peak}), and several noticeable maxima in the region of $A\\gtrsim 100$, for the distributions of both odd- and even-A species.  \\begin{figure} \\includegraphics[scale=.70]{lodjj.eps} \\caption{Solar System isotopic abundances versus mass number, $A$. Abundances are based on meteoritic samples as well as on values inferred from the solar photosphere (Sect. 2). Values taken from Lodders \\cite{Lod03}. \\label{SSabun}} \\end{figure}  The reasons for favoring the presence of some species and for the specific pattern inferred for the Solar System abundances have spawned large debates. Following earlier attempts, Hoyle proposed in 1946 that all heavy elements are synthesized in stars \\cite{Hoy46}. A major breakthrough was achieved with the detection of technetium, an element with no stable isotopes, in the spectra of several S stars by Merrill \\cite{Mer52} in the early 1950s. Since the technetium isotopes have rather short half-lives on a Galactic time scale, the obvious conclusion was that technetium must have been produced in the observed stars and thus nucleosynthesis must still be going on in the Universe. Shortly after, two seminal papers that provided the theoretical framework for the  origin of the {\\it chemical} species elements, were published (almost exactly a century after Darwin's treatise on the origin of {\\it biological} species) by Burbidge, Burbidge, Fowler \\& Hoyle \\cite{Bur57}, and independently, by Cameron \\cite{Cam57}. These works identified stars as Galactic factories that synthesize atomic nuclei via nuclear reactions. The stars return this processed material to the interstellar medium through various hydrostatic or explosive means during their evolution and thereby enrich the interstellar medium in {\\it metals}, that is, any species heavier than He according to the astronomical jargon. Out of this matter new generations of stars are born and a new cycle of nuclear transmutations will be initiated. In fact, without concourse of the many nuclear processes that take place in the stars and that are discussed in this review, the Universe would have remained a chemically poor place, and probably no life form would have ever emerged.  This review summarizes more than 50 years of progress in our understanding of the nuclear processes that shaped the current chemical abundance pattern of the Universe. The structure of the paper is as follows. In Section 2, we briefly address the main aspects associated with the determination of the Solar System abundances. This is followed by an overview on Big Bang nucleosynthesis (Sect. 3), and by an account of the nuclear processes governing low-mass stars (Sect. 4). Particular attention is paid to the Standard Solar model and to lessons learned from neutrino physics. In Sect. 5, we outline the basic properties of red giant and AGB stars, with emphasis on the {\\it s-process}. The series of explosive processes that characterize core-collapse supernovae and $\\gamma$-ray bursts (e.g., {\\it p-, $\\nu$-, $\\nu$p-, $\\alpha$-, and r-processes}) are reviewed in Sect. 6. This is supplemented with an overview of different astrophysical scenarios that involve binary stellar systems, such as classical and recurrent novae (Sect. 7), type Ia supernovae (Sect. 8), X-ray bursts and superbursts (Sect. 9), and stellar mergers (Sect. 10). We briefly address as well non-stellar processes, such as spallation reactions (Sect. 11), and conclude with a number of key puzzles and future perspectives in this field (Sect. 12). ",
        "conclusions": "The last fifty years have witnessed extraordinary progress in the modeling of stellar evolution, as well as in our understanding of the origin of the elements. First, supercomputers have allowed astrophysicists to perform complex hydrodynamic simulations coupled with extensive nuclear reaction networks at astounding levels of resolution. Indeed, last decades have seen the emergence of multidimensional simulations in many areas of stellar astrophysics research. Second, better observational constraints have been obtained through combined efforts in spectroscopy, the dawn of high-energy astrophysics with space-borne observatories, and laboratory measurements of presolar meteoritic grains that condensed in a suite of stellar environments. And third, nuclear physics has witnessed the birth of radioactive ion beam facilities, as well as direct measurements at astrophysically important energies using stable ion beam facilities.  Further progress eagerly awaits future developments aimed at solving key remaining questions in nuclear astrophysics, chief among them: \\begin{enumerate} \\item Why do predictions of helioseismology disagree with those of the standard solar model? \\item What is the solution to the lithium problem in Big Bang nucleosynthesis? \\item What do the observed light-nuclide and s-process abundances tell us about convection and dredge-up in massive stars and AGB stars? \\item What are the production sites of the $\\gamma$-ray emitting radioisotopes $^{26}$Al, $^{44}$Ti and $^{60}$Fe? \\item What is the origin of about 30 rare and neutron deficient nuclides beyond the iron peak (p-nuclides)? \\item What causes core-collapse supernovae to explode? \\item What is the extend of neutrino-induced nucleosynthesis ($\\nu$-process)? \\item What is the extend of the nucleosynthesis in proton-rich outflows in the early ejecta of core-collapse supernovae ($\\nu p$-process)? \\item What are the sites of the r-process? \\item What causes the discrepancy between models and observations regarding the mass ejected during classical nova outbursts? \\item Which are the physical mechanisms driving convective mixing in novae? \\item What are the progenitors of type Ia supernovae? \\item What is the nucleosynthesis endpoint in type I X-ray bursts? Is there any matter ejected from those systems? \\item What is the impact of stellar mergers on Galactic chemical abundances? \\item What are the production and acceleration sites of Galactic cosmic rays? \\end{enumerate}  \\noindent These questions will challenge young and courageous researchers for years to come, thereby keeping nuclear astrophysics in the spot light of science."
    },
    "1107/1107.2849_arXiv.txt": {
        "abstract": "{ The role of disks in the formation of high-mass stars is still a matter of debate but the detection of circumstellar disks around O-type stars would have a profound impact on high-mass star formation theories. }{ We made a detailed observational analysis of a well known hot molecular core lying in the high-mass star-forming region G31.41+0.31. This core is believed to contain deeply embedded massive stars and presents a velocity gradient that has been interpreted either as rotation or as expansion, depending on the authors. Our aim was to shed light on this question and possibly prepare the ground for higher resolution ALMA observations which could directly detect circumstellar disks around the embedded massive stars. }{ Observations at sub-arcsecond resolution were performed with the Submillimeter Array in methyl cyanide, a typical hot molecular core tracer, and \\CO\\ and \\COI, well known outflow tracers. We also obtained sensitive continuum maps at 1.3~mm. }{ Our findings confirm the existence of a sharp velocity gradient across the core, but cannot confirm the existence of a bipolar outflow perpendicular to it. The improved angular resolution and sampling of the uv plane allow us to attain higher quality channel maps of the \\MCN\\ lines with respect to previous studies and thus significantly improve our knowledge of the structure and kinematics of the hot molecular core. }{ While no conclusive argument can rule out any of the two interpretations (rotation or expansion) proposed to explain the velocity gradient observed in the core, in our opinion the observational evidence collected so far indicates the rotating toroid as the most likely scenario. The outflow hypothesis appears less plausible, because the dynamical time scale is too short compared to that needed to form species such as \\MCN, and the mass loss and momentum rates estimated from our measurements appear too high. } ",
        "introduction": "\\label{sint}  High-mass stars are usually defined as those exceeding $\\sim$8~\\Msun, based on the fact that stars above this mass limit do not have a pre-main sequence phase (Palla \\& Stahler \\cite{past}). This means that accretion is ongoing until the star ignites hydrogen burning and reaches the zero-age main sequence (ZAMS). At this point the strong radiation pressure may halt and even reverse infall and thus stop further growth of the stellar mass. This led to the so-called ``radiation pressure problem''. Recent studies have demonstrated that this limitation holds only in spherical symmetry. As first envisioned by Nakano~(\\cite{nakano}) and recently demonstrated by Krumholz et al.~(\\cite{krum}) and Kuiper et al.~(\\cite{kuip}), accretion through a circumstellar disk can explain the formation of stars up to the upper limit of the initial mass function, by allowing part of the photons to escape along the disk axis and boosting the ram pressure of the accreting gas through the small disk solid angle. It also appears that the powerful ionizing fluxes from these OB-type stars are not sufficient to destroy the disk, which eventually turns into an ionized, rotating accretion flow close to the star (Sollins et al.~\\cite{sollins}; Keto \\cite{keto07}).  For these reasons it seems established that circumstellar accretion disks play a crucial role in the formation of {\\it all} stars and not only solar-type stars. This theoretical result contrasts with the limited observational evidence of disks in high-mass (proto)stars. Only in recent years the number of disk candidates associated with luminous young stellar objects (YSOs) has significantly increased, mostly owing to the improvement of (sub)millimeter interferometers in terms of angular resolution and sensitivity. The main problems that one has to face in this type of search are the large distances to the sources (typically a few kpc) and the confusion caused by stellar crowding (OB-type stars form in clusters). These factors may explain the failure to detect disks in O-type stars, as opposed to the number of detections obtained for B-type stars (Cesaroni et al.~\\cite{ppv}). In association with the most luminous YSOs one finds only huge ($\\la$0.1~pc), massive (a few 100~\\Msun) cores, with velocity gradients suggesting rotation. These objects, named ``toroids'', are likely non-equilibrium structures, because the ratio between the accretion time scale and the rotation period is shorter than for disks: this implies that the toroid does not have enough time to adjust its structure to the new fresh material falling onto it (Cesaroni et al.~\\cite{ppv}; Beltr\\'an et al.~\\cite{bel11}).  With this in mind, one can see that understanding the formation of the most massive stars ($>20~\\Msun$) may benefit from a detailed investigation of toroids, also because these objects might be hosting true circumstellar disks in their interiors. Moreover, their mere existence may set tighter constraints on theoretical models. Despite the number of candidates, the existence of rotating toroids is still a matter of debate. What is questioned is the nature of the velocity gradient, which is sometimes interpreted as expansion instead of rotation (see e.g. Gibb et al.~\\cite{gibb} and Araya et al.~\\cite{araya} and references therein), thus suggesting that one might be seeing a compact bipolar outflow rather than a rotating core. In an attempt to distinguish between these two possibilities and more in general to shed light on these intriguing objects, we have focused our attention on one of the best examples: the hot molecular core (HMC) G31.41+0.31 (hereafter G31.41).  This prototypical HMC is located at a kinematic distance of 7.9~kpc and was originally imaged in the high-excitation (4,4) inversion transition of ammonia (Cesaroni et al.~\\cite{cesa94}) and in the (6--5) rotational transitions of methyl cyanide (\\MCN; Cesaroni et al.~\\cite{cesag31}). The latter showed the existence of a striking velocity gradient (centered at an LSR velocity of $\\sim$96.5~\\kms) across the core in the NE--SW direction, already suggested by the distribution and velocities of OH masers (Gaume \\& Mutel~\\cite{gamu}). Follow-up interferometric observations with better angular resolution and in high-energy tracers have confirmed this result and revealed the presence of deeply embedded YSOs, which in all likelihood explains the temperature increase toward the core center (Beltr\\'an et al.~\\cite{bel04}, \\cite{bel05}, hereafter BEL04 and BEL05; Cesaroni et al.~\\cite{cesa10}). The G31.41 HMC is separated by $\\sim$5\\arcsec\\ from an ultracompact (UC) \\HII\\ region, and overlaps in projection on a diffuse halo of free-free emission, possibly associated with the UC~\\HII\\ region itself (see e.g. Fig.~2c of Cesaroni et al.~\\cite{cesa98}). The SPITZER/GLIMPSE images (Benjamin et al.~\\cite{benj}) show that the HMC lies in a complex pc-scale region where both extended emission and multiple stellar sources are detected. All these facts complicate the interpretation of the HMC structure and its relationship with the molecular surroundings and call for additional high-quality observations.  An important aspect of a proper interpretation of the velocity gradient observed in the HMC is the existence of an associated molecular outflow. A reasonable assumption, suggested by the analogy with low-mass YSOs, is that a disk actively undergoing accretion must be associated with a bipolar jet/outflow expanding along the disk rotation axis. Therefore, a way to distinguish between the two interpretations of the HMC velocity gradient (expansion vs rotation) is to search for a molecular outflow perpendicular to the gradient. Previous interferometric observations in the \\COI(1--0) line have suggested the presence of a collimated outflow directed SE--NW (Olmi et al.~\\cite{olmi96b}), but the poor uv coverage, especially on the shortest baselines, makes this result questionable. We have therefore conducted new interferometric observations at 1.3~mm with the Submillimeter Array (SMA) in the typical HMC tracer \\MCN\\ and in standard outflow tracers such as \\CO\\ and \\COI. In order to be sensitive to extended structures filtered out by the interferometer, we have also mapped the region over $\\sim$2\\arcmin\\ with the IRAM 30-m telescope. The observational details are given in Sect.~\\ref{sobs}, while the results are illustrated in Sect.~\\ref{sres} and discussed in Sect.~\\ref{sdis}. Finally, the conclusions are drawn in Sect.~\\ref{scon}. ",
        "conclusions": "\\label{scon}  With the present study we have performed a deep analysis of the geometrical structure and kinematics of the HMC in the high-mass star-forming region G31.41. Thanks to new high angular resolution SMA images and complementary IRAM 30-m maps, we have investigated both the HMC, through the \\MCN(12--11) transition, and the lower density surroundings, through \\CO\\ and \\COI. We could thus search for a possible outflow powered by the star(s) embedded in the HMC. In particular, our main goal was to shed light on the nature of the NE--SW velocity gradient, previously detected in the HMC by various authors, and possibly distinguish between rotation and expansion.  Our new data confirm the presence of the velocity gradient basically in all HMC tracers observed, and indicate that a red--blue symmetry along the same axis is seen also in the CO isotpomers. We were unable to detect any bipolar outflow along the SE--NW direction, in contrast with the findings of Olmi et al.~(\\cite{olmi96b}). The latter was probably an artifact caused by very limited uv sampling and lack of zero-spacing information.  We find a hint of infall in the HMC, consistent with the results of Girart et al.~(\\cite{gira}). A comparison between the \\MCN\\ and CO maps suggests that these species are tracing the same phenomenon. Especially the \\COI(2--1) line emission appears to be the obvious prosecution of the \\MCN(12--11) emission on a larger scale.  We have discussed whether the velocity gradient observed in G31.41 is a compact bipolar outflow or a rotating, geometrically thick, toroidal structure. The former hypothesis implies outflow parameters typical of high-mass stars with luminosities in excess of $10^5$~\\Lsun, whereas previous estimates by Osorio et al.~(\\cite{osor}) and Cesaroni et al.~(\\cite{cesa94,cesa10}) indicate a significantly lower luminosity for G31.41. Moreover, such an outflow would be much more compact than typical bipolar outflows detected in high-mass star-forming regions and the dynamical time scale would be an order of magnitude shorter than that needed to form HMC spacies such as \\MCN. It is also worth noting that a core undergoing expansion seems difficult to reconcile with the detection of infall.  The composite \\MCN\\ and \\COI\\ position-velocity plot can be explained if the gas is undergoing both (pseudo-)Keplerian rotation and infall, suggesting that this massive object could be dynamically stabilized by a compact cluster of YSOs tightly packed inside a few 1000~AU from the center. In this scenario the \\CO\\ emission could be tracing the outer regions, where the gravitational field is dominated by the gas and the rotation velocity tends to a constant value. Albeit intriguing, this interpretation remains speculative and could be proved only by producing synthetic \\MCN\\ and CO maps and position-velocity plots to be compared with the data. This numerical effort goes beyond the purposes of the present study.  We conclude that the case of G31.41 could be a scaled-up version of lower-mass YSOs, such as GG~Tau and MWC\\,758, where circumstellar disks have been detected and successfully modeled. We stress that whatever the interpretation, the observations of this object have produced an important finding for the study of high-mass star formation. We obtained iron-clad evidence for the existence of a velocity gradient across the HMC and demonstrated that this gradient is {\\it not} due to multiple cores unresolved in the beam and moving at different velocities, but to the gas undergoing a smooth drift in (line-of-sight) velocity from the one end to the other of the core. If, as we believe, one is dealing with a rotating toroid, then this indicates that star formation in this HMC may proceed in a similar way to low-mass stars, namely through circumstellar, centrifugally supported disks ``hidden'' inside the densest interiors of the toroid itself. If, instead, the outflow interpretation were correct, such a sharp, neat velocity gradient would strongly suggest the existence of a very effective focusing mechanism for the outflow and the best candidate as a collimating agent would be a (yet undetected) circumstellar disk. Therefore, regardless of whether the G31.41 HMC is undergoing expansion or rotation, the role of disks in the mechanism of high-mass star formation seems to be strengthened by our findings. The advent of ALMA will be crucial to detect deeply embedded disks in HMCs."
    },
    "1107/1107.5658_arXiv.txt": {
        "abstract": "Ultra-high energy charged particles of unknown origin, which interact in the high atmosphere of the Earth generating large cascades of secondary particles, can be observed from the ground. For many decades, they have been a puzzle for particle physicists and astrophysicists. They seem to arrive from random directions in the sky, although the most energetic ones are supposed to point towards their sources with good accuracy.  As an attempt to discriminate among several possible production scenarios, astrophysicists try to test the statistical isotropy of the directions of arrival of these cosmic rays.  At the highest energies however, the observed cosmic rays are very rare, and testing the distribution of such small samples of directional data on the sphere is non trivial. We address here a nonparametric detection problem, making use of a multiscale analysis of the observation sample, based on a decomposition of the directional data using a wavelet frame on the sphere, the needlets. The aim is to test the isotropy, or the equality of the distribution with a known one.  We explore two particular procedures, a multiple test and a plug-in approach.  We describe the practical implementation of these two procedures and compare them to other methods in the literature. As alternatives to isotropy, we consider both very simple toy-models, and more realistic non isotropic models based on Physics-inspired simulations. The Monte Carlo study shows the good performances of the multiple test at moderate sample size, together with the robustness of its sensitivity with respect to the unknown characteristics of the alternative hypothesis.  On the 69 most energetic events published by the Pierre Auger collaboration, our procedure detects significant departure from isotropy. The flexibility of this method and the possibility to modify it to take into account a large variety of extensions of the problem make it an interesting option for future investigation of the origin of ultra-high energy cosmic rays. ",
        "introduction": "\\label{sec:introduction}  \\subsection{Motivations} \\label{sec:motivations}   It is a common problem in astrophysics to analyse data sets containing measurements of a number of objects (such as galaxies of a particular type) or of events (such as cosmic rays or gamma ray bursts) distributed on the celestial sphere. Each set of such objects or events can be represented as a collection of positions $X_i = (\\theta_i,\\phi_i), i=1,\\dots,n$ in $\\Sset$ the unit sphere of $\\Rset^3$. In many cases, such objects trace an underlying probability distribution $f$ on the sphere, which itself depends on the physics which governs the production of the objects and events. Galaxies, for instance, form in over-densities of a preexisting smooth field of distribution of matter in the universe, and the study of the statistics of their distribution has grown into a field of astrophysics by itself \\citep{2002sgd..book.....M}.  The case of ultra high energy cosmic rays (UHECRs) is of particular interest, and is the main focus of the present work. UHECRs are particles of unknown origin which arrive at the Earth from apparently random directions of the sky. These particles interact with atoms of the upper atmosphere, generating a huge cascade of billions of secondary particles. The observation of these secondary particles with appropriate detectors on ground permits the measurement of the direction of arrival and of the energy of the original cosmic ray.  The existence of cosmic rays has been known for about a century. Such particles exist with a very wide range of kinetic energies, from few eV to more than $10^{20}$ eV.\\footnote{1 eV = 1 electron Volt $\\simeq 1.6 \\times 10^{-19}$ Joule} Observed cosmic rays are typically ordinary charged particles (electrons, protons and nuclei), propagating in empty space, and deflected by galactic magnetic fields. The rate of observed cosmic rays in the vicinity of the Earth, however, decreases rapidly with energy. At low energy, the observed cosmic rays are numerous and their composition is well known. There also exist several known astrophysical processes responsible for their acceleration, such as stellar winds for the least energetic ones, to violent phenomena such as supernovae shock waves at higher energy. At the highest energies ($E \\geq 10^{20}$ eV), however, the observed flux is of the order of 1 event per square kilometre  per century, which limits the statistics of observed events to few tens of events (in two decades of observations). In addition, no understood astrophysical process, involving known objects, can accelerate particles to such tremendous energies.  Recent observations of ultrahigh-energy cosmic rays suggest that they are ordinary particles, such as protons and nuclei, accelerated in extremely violent astrophysical phenomena \\citep[see][for a recent review on the astrophysics of UHECRs]{2011ARA&A..49..119K}. However, many alternate hypotheses concerning their nature and origin have been proposed over the years \\citep[see, e.g.,][]{1984ARA&A..22..425H,2004RPPh...67.1663T,2005NuPhS.138..465C}. UHECRs could originate from active galactic nuclei (AGN), or from neutron stars surrounded by extremely high magnetic fields, or yet from many other processes. It is also possible that the type and origin of ultra high energy cosmic rays (at energies above $10^{19}$ eV) depend, at least to some extent, upon the energy at which they are observed. Indeed, the most energetic cosmic rays cannot propagate very far (i.e. not much more than $\\sim 100 $  Mpc), without loosing most of their energy by interactions with photons from the Cosmic Microwave Background \\citep[the so-called GZK effect;][]{1966PhRvL..16..748G,1966JETPL...4...78Z}. The confirmation of the energy cut-off at the high end of the cosmic ray spectrum is one of the main achievements of the Pierre Auger Observatory \\citep{2008PhRvL.101f1101A,2010PhLB..685..239A}.  Before the location and physical process of acceleration have been clearly  identified, taking into account the fact that most of the evidence about the chemical composition of  cosmic rays at the highest energies rely on extrapolations of the present knowledge of hadronic interactions at energies two orders of magnitude above the range presently tested at the LHC, it is difficult to completely rule-out  alternate theoretical explanations as to what UHECRs exactly are and what is their origin. Alternate hypotheses such as production by decay of long-lived relic particles from the Big Bang, about 13 billion years old \\citep{2000PhR...327..109B}, are just starting to be disfavored by the observations of the Auger collaboration, with recently published results about primary photon limits that impose stringent limits on these kind of models \\citep{2009APh....31..399P}.  In an attempt to better understand the origin of such UHECRs, physicists study the statistical distribution of their directions of arrival, looking for two particular signatures. First, the (statistically significant) arrival of more than one UHECR from the same direction on the sky would indicate that their production is not likely to originate from single time events (e.g. catastrophic mergers of two compact astrophysical objects), but rather from sources which emit UHECRs regularly.\\footnote{With the caveat that the time of propagation may depend on the energy and on the exact trajectory followed by the UHECR to reach us, making it possible that two particles reaching the Earth at different times have actually been emitted simultaneously.} Second, one may look for correlation in the directions of arrival of UHECRs with known astrophysical objects, as nearby active galactic nuclei, in an attempt to identify plausible production sites. Hence, in some hypotheses, the underlying probability distribution for the directions of incidences of observed UHECRs would be a finite sum of point-like sources -- or nearly point-like, taking into account the deflection of the cosmic rays by magnetic fields. In other hypotheses, the distribution could be uniform, or smooth and correlated with the local distribution of matter in the universe. The distribution could also be a superposition of the above. Distinguishing between these hypotheses is of primordial importance for understanding the origin and mechanism of production of UHECRs.  In the past 20 years, a number of experiments have gathered observations of UHECRs, and several papers have been written which look for such features in the distribution of their directions of arrival, with sometimes contradictory conclusions. The difficulty lies in the fact that UHECRs are rare, and that they do not arrive necessarily exactly from the direction where their source is located. Indeed, as typical cosmic ray particles are charged (which permits their acceleration by electromagnetic processes), they are deflected by Galactic and intergalactic magnetic fields. The deflection depends on the length of the path through the magnetic field, and on the energy and charge of the particle. In fact, only very energetic cosmic rays (above few $10^{19}$ eV) with small charge (e.g. protons or nuclei with small atomic numbers) are expected to travel typical astrophysical distances from their source to us with deflection angles smaller than a few degrees. Details of the deflections are not known, as neither the exact magnitude, orientation, and regularity on large scales of Galactic and extragalactic magnetic fields, nor the distance of the sources of UHECRs, nor the exact energy of the incoming cosmic ray, nor its charge (to within a factor of 26 between protons and iron nuclei), are known. Errors on the direction of the source of an UHECR can then be of order $1^\\circ$ at the lowest (typical error on the measurement of the direction of arrival with Auger), up to few degrees for protons, or tens of degrees for heavy nuclei travelling a long path through a regular galactic magnetic field.  Given a set of observed UHECRs, how can one best test for ``repeaters'' (cosmic rays coming from the same source) or more generally anisotropy in the distribution? If one restricts the analysis to the few events for which one is sure that the deflection angle is negligible, events are scarce and there is not enough statistics to conclude. As one selects events with less energy, the direction of origin becomes less reliable, with the total number of events completely dominated by those events with poorly constrained direction of origin. Finally, it is not clear how to build the isotropy test, without any sound prior knowledge about the uncertainty in the measured direction of the source. All of these are very meaningful questions to analyse UHECR observations.  Recently, an analysis of the direction of arrival of 27 UHECRs observed by the Pierre Auger experiment concludes in the existence of an anisotropy, and a correlation with objects in a catalogue of nearby active galactic nuclei (AGNs), located at distances lower than about 70 Mpc\\footnote{70 million parsecs $\\simeq 2.15 \\times 10^{21}$ km} \\citep{2008PhRvL.101f1101A}. This anisotropy, however, is less obvious in a more recent analysis, based on 69 observed events \\citep{2010APh....34..314T}. Clearly, the statistics is limited, and the development of new methods for investigating this topic can provide new insights on the origin of the UHECRs. Methods independent of external data sets such as the fore-mentioned VCV catalogue (which is not a statistically well-characterized sample of AGNs but a compilation of published results) are of particular interest.  \\subsection{Outline of this work} \\label{sec:outline-paper}   This work focuses on the important question of the isotropy of the cosmic rays. Because of the small number of available data, this question is not answered yet, although data from the Pierre Auger collaboration seems to hint at a correlation between the directions to the ultra-high energetic events (above $5.5 \\times 10^{19}$ eV) and the directions to active galactic nuclei in the catalogue compiled by  V\\'eron \\& Cetty-V\\'eron   \\citep[see][]{2008APh....29..188P,2010APh....34..314T}.  From a statistical point of view, we address the question of testing the goodness-of-fit of the isotropy assumption to this small sample of directional data. The framework we choose is purely nonparametric as we do not want to favour any particular alternative hypothesis, and as we wish to be able to discover unexpected forms of anisotropy.  The paper is organized as follows. In Section~\\ref{toymodel}, we present a simplified model of cosmic ray propagation which will be used in Monte Carlo simulations to test the method. In Section \\ref{sec:nonp-tests-anis} we present the nonparametric framework. Then we describe our needlet based anisotropy tests in Section~\\ref{sec:anisotropy-tests}.  In section~\\ref{sec:monte-carlo-exper}, we present a Monte Carlo experiment that compares the power of the different test in terms of sensitivity and robustness with respect to the parameters of the methods. We apply our procedures to real data from the Pierre Auger collaboration in Section~\\ref{sec:analysis-auger-data}. We then conclude and give perspectives for future extensions of the present work. Appendix~\\ref{sec:needl-constr-pract} in the on-line supplemennt is devoted to a short description of the type of wavelets we have used (the needlets) and the practical and numerical implementation of our methods. ",
        "conclusions": "In this paper, we have investigated the problem of the detection of anisotropy of directional data on the unit sphere, with an application to the analysis of ultra high energy cosmic ray events as observed with a detector such as the Pierre Auger observatory. It was important to consider samples whose sizes are comparable to the sizes of the data sets that are available nowadays for cosmic rays scientists (about 25 at the beginning of this work, about a hundred now). Although we are mainly interested in small sample performances, we  have proposed a multiple test approach based on a multiresolution analysis of the data, which could hopefully be proved to be  asymptotically optimal in the minimax sense, a well-known pessimistic framework.  We have proposed, and tested on various simulated data sets, two methods using the decomposition of the directional data onto a frame of spherical needlets. Their performance has been compared to other (more specific) approaches based on the nearest neighbour and on the two point correlation function. The simulation shows that the needlet-based methods perform comparatively very well in various situations. They are competitive with existing method at small sample size, and tend to outperforms them from moderate sample size.  Moreover, the ``omnibus'' property of the needlets method is interesting for the problem at hand, in which the type of possible anisotropy (the class of alternative) is not really well known a priori. In addition, a multiple test based on the use of spherical needlets offers a good opportunity to extend the method of detection of anisotropies with not only multiplicity in the scales tested, but also in ranges of energy of the incoming particles. Indeed, while in this work we have used the energy level as a simple threshold, one could instead implement a detection using the joint directional-energy information -- allowing thus to simultaneously extract information from the highest energy cosmic rays, which are not deflected much by Galactic and extragalactic magnetic fields, and also from lower energy events, more deflected but much more numerous. In the light of our simulations on an energy level-dependent model, multiscale approach could lead to stronger conclusion using the CR data that are not yet made public by the Pierre Auger Observatory.  As in any nonparametric method, there is at least one parameter to be tuned, often by hand or using more sophisticated data-driven methods such as cross-validation. In the needlet methods one can tune several parameters (shape of the needlets, highest scale \\bandmax --- although there's is an asymptotic formula for it---, thresholds on the coefficients in the \\plugin{} approaches, thresholds on the individual tests in the \\multiple{} procedure, power norm). It is plausible however that a large range of possible choices for most of these parameters give comparable performance.  Although we have used needlets that are compactly supported windows in the harmonic space, it may be arguable that they are not the most appropriate tool. One could consider, as an alternative, better spatially concentrated functions \\citep[such as the Mexican needlets, see \\eg][]{lan:marinucci:2009} or, in general, try to optimize the needlet window function given prior knowledge of the physical problem and of the expected properties of anisotropic distributions of the cosmic ray direction of incidence. In this spirit, it would be interesting to consider directional wavelet such as curvelets or ridgelets \\citep[see][]{starck+2006} to test for specific strip-like alternative densities.  It is also possible to consider non-dyadic needlets. The choice of $\\cB \\in (1,2)$ allows a finer coverage of the frequency line. The numerical results presented here have not taken this benefit into full account, and whether significantly higher power can be obtained by optimizing this number remains to be investigated.  Finally, in addition to the aforementioned possible extensions of our methods, we want to stress that the work presented here also opens the way to two lines of future investigations, one on the applications side, and one more theoretical. On the experimental side, it will be of much interest to apply the method on larger data sets (for instance by lowering the energy threshold to increase the available sample size). On the theoretical side, the validation of the approach has to be investigated on the basis of some theory in the minimax framework it is designed for.   \\subsubsection*{Acknowledgment} \\label{sec:acknowledgment}  For the numerical part of this work, we acknowledge the use of the \\textsc{Healp}ix package \\citep{gorski+2005} and of the SphereLab (an \\textsc{Octave} interface to \\textsc{Healp}ix package, needlet transforms, and utilities).  We thank Dmitri Semikoz for useful discussions concerning high energy cosmic rays, and Maude le Jeune, Jean-Fran\u00e7ois Cardoso and Fr\u00e9d\u00e9ric Guilloux for making available some of the tools used for the computational aspects of the present work. We thank the anonymous referees and associate editor for their suggestions and remarks that improved the presentation of this paper.     \\tiny"
    },
    "1107/1107.4096.txt": {
        "abstract": "We estimate the bulk Lorentz factor \\G\\ of  31 GRBs using the measured peak time of their afterglow light curves. We consider two possible scenarios for the estimate of \\G: the case of a homogeneous circumburst medium or a wind density profile. The values of \\G\\ are broadly distributed between few tens and several hundreds with average values $\\sim$138 and $\\sim$66 for the homogeneous and wind density profile, respectively. We find that the isotropic energy and luminosity correlate in a similar way with \\G, i.e. \\eiso$\\propto$\\G$^2$ and \\liso$\\propto$\\G$^2$, while the peak energy \\ep$\\propto$\\G. These correlations are less scattered in the wind density profile than in the homogeneous case. We then study the energetics, luminosities and spectral properties of our bursts in their comoving frame. The distribution of \\lisocom\\ is very narrow with a dispersion of less than a decade in the wind case, clustering around  \\lisocom$\\sim5\\times10^{48}$ erg s$^{-1}$. Peak photon energies cluster around \\epcom$\\sim$ 6 keV. The newly found correlations involving \\G\\ offer a general interpretation scheme for the spectral--energy correlation of GRBs. The \\ama\\ and \\yone\\ correlations are due to the different \\G\\ factors and the collimation--corrected correlation, \\ghi\\ (obtained by correcting the isotropic quantities for the jet opening angle $\\theta_{\\rm j}$), can be explained if $\\theta_{\\rm j}^2\\Gamma_{0}=$ constant. Assuming the \\ghi\\ correlation as valid, we find a typical value of $\\theta_{\\rm j}$\\G$\\sim$  6--20, in agreement with the predictions of magnetically accelerated jet models. ",
        "introduction": "The discovery of the afterglows of Gamma Ray Bursts (GRBs - Costa et al. 1997) allowed to pinpoint their position in the X--ray and Optical bands. This opened a new era focused at measuring the spectroscopic redshifts of these sources. The present\\footnote{http://www.mpe.mpg.de/$\\sim$jcg/grbgen.html} collection of GRBs with measured $z$ consists of 232 events. In 132 bursts of this sample (updated in this paper) the peak energy \\epo\\ of their $\\nu F_{\\nu}$ prompt emission $\\gamma$--ray spectrum could be constrained. In turn, for these bursts it was possible to calculate the isotropic equivalent energy \\eiso\\ and luminosity \\liso. The knowledge of the redshifts showed that two strong correlations exist between the {\\it rest frame} peak energy \\ep\\ and \\eiso\\ or \\liso\\ (also known as the ''Amati'' and ''Yonetoku\" correlations -- Amati et al. 2002, Yonetoku et al. 2004, respectively).  The reality of these correlations has been widely discussed in the literature. Some authors pointed out that they can be the result of observational selection effects (Nakar \\& Piran 2005; Band \\& Preece 2005; Butler et al. 2007, Butler, Kocevski \\& Bloom 2009; Shahmoradi \\& Nemiroff 2011) but counter--arguments have been put forward arguing that selection effects, even if surely present, play a marginal role (Ghirlanda et al. 2005, Bosnjak et al. 2008, Ghirlanda et al. 2008; Nava et al., 2008; Krimm et al. 2009; Amati et al. 2009). The finding that a correlation $E_{\\rm p}(t)$--$L_{\\rm iso}(t)$ exists when studying time--resolved spectra of individual bursts is a strong argument in favor of the reality of the spectral energy correlations, (Ghirlanda, Nava\\& Ghisellini 2010; Ghirlanda et al. 2011) and motivates the search for the underlying process generating them. Even if several ideas have been already discussed in the literature, there is no general consensus yet, and a step forward towards a better understanding both of the spectral energy correlations and the underlying radiation process of the prompt emission of GRBs is to discover what are the typical energetics, peak frequencies and peak luminosities in the {\\it comoving frame}.  The physical model of GRBs requires that the plasma emitting $\\gamma$--rays should be moving relativistically with a bulk Lorentz factor \\G\\ much larger than unity. The high photon densities and the short timescale variability of the prompt emission imply that GRBs are optically thick to pair production which, in turn, would lead to a strong suppression of the emitted flux, contrary to what observed. The solution of this compactness problem requires that GRBs are relativistic sources. From this argument lower limits $\\Gamma_{0} \\ge 100$ are usually derived (Lithwick \\& Sari, 2001). The first observational evidences supporting this scenario were found in the radio band where the ceasing of the radio flux scintillation (few weeks after the explosion as in GRB 970508; Frail et al. 1997), allowed to estimate \\GA\\ of a few. This value corresponds to the late afterglow phase, when the fireball is decelerated almost completely by the interstellar medium and is characterized by a much smaller bulk Lorentz factor than the typical \\G\\ of the prompt phase.  Large Lorentz factors imply strong beaming of the radiation we see. We are used to consider GRB intrinsic properties (\\ep, \\eiso, \\liso) for the bursts with measured redshifts, but still an important correction should be applied. Our aim is to study the distributions of \\ep, \\eiso, \\liso\\ and the spectral--energy correlations (\\ama\\ and \\yone) in the {\\it comoving frame}, accounting for the \\G\\ factor. The estimate of \\G\\ is possible by measuring the peak of the afterglow (Sari \\& Piran 1999) and has been successfully applied in some cases (e.g. Molinari et al. 2007, Gruber et al., 2011) and more extensively recently by Liang et al. (2010) in the optical and X--ray band. Other methods allow to set lower limits (Abdo et al. 2009; Ackerman et al. 2010; Abdo et al. 2009a) mainly by applying the compactness argument to the high energy emission recently detected in few GRBs at GeV energies by the \\fe\\ satellite (see Zou, Fan \\& Piran 2011; Zhao, Li \\& Bai 2011; Hascoet et al. 2011 for more updated calculation on these lower limits on \\G). Conversely, upper limits (Zou \\& Piran 2010) can be derived by requiring that the forward shock emission of the afterglow does not appear in the MeV energy band.  The paper is organized as follows: in \\S~2 we discuss the relativistic corrections that allow us to derive the comoving frame \\epcom, \\eisocom and \\lisocom\\ from the rest frame \\ep, \\eiso, \\liso; in \\S~3 and \\S~4 we derive a general formula for the estimate of \\G\\ from the measurement of the time of the peak of the afterglow emission; in \\S~5 we present our sample of GRBs and in \\S~6 our results which are finally discussed in \\S~7. Throughout the paper we assume a standard cosmology with $h=\\Omega_{\\Lambda}=0.7$ and $\\Omega_{m}=0.3$. ",
        "conclusions": "We have considered all bursts with measured $E_{\\rm peak}$ and known redshift up to May 2011 (132 GRBs). Among these we have searched in the literature for any indication of the peak of the afterglow light curve $t_{\\rm p,z}$ suitable to estimate the initial bulk Lorentz factor \\G. Our sample of bursts is composed by  27 GRBs with a clear evidence of $t_{\\rm p,z}$  in the optical light curve. We have derived the peak energy $E^\\prime_{\\rm peak}$, the isotropic energy $E^\\prime_{\\rm iso}$ and the isotropic peak luminosity $L^\\prime_{\\rm iso}$ in the comoving frame. To this aim we have derived the general formula for the computation of \\G\\ (\\S.3) considering two possible scenarios: a uniform interstellar medium density profile ($n=$const, H) or a wind density profile ($n \\propto r^{-2}$, W).  For the wind case the \\G\\--distribution (Fig. \\ref{fg1} and Tab. \\ref{tab2}) is shifted at somewhat smaller values ($\\langle\\Gamma_{0}\\rangle\\sim$ 66) than the same distribution for the homogeneous density case ($\\langle\\Gamma_{0}\\rangle\\sim$ 138). The distribution of $E^\\prime_{\\rm peak}$ is relatively narrow and centered around $\\sim$6 keV or $\\sim 3$ keV for the W and H case (Fig. \\ref{fg3} and Tab. \\ref{tab2}). The distribution of $L^\\prime_{\\rm iso}$ (Fig. \\ref{fg5}) clusters, especially for the wind case, in a very narrow range (much less than a decade), around $5\\times10^{48}$ erg s$^{-1}$, while the distribution of $E^\\prime_{\\rm iso}$ (Fig. \\ref{fg4}) is broader  and centered at $3\\times10^{51}$ erg. $E_{\\rm iso}$ and $L_{\\rm iso}$ correlate with \\G, ($\\propto$\\G$^{2.2}$ both for the wind and the homogeneous case) and the correlation is stronger (with a scatter $\\sigma=0.07$) for the wind case. Finally, the duration of the burst, as expected, does not correlate with \\G.  The correlations that we have found are strong despite they are defined with a still small number of GRBs. We expect that with the increase of the number of GRBs with measured $t_{p,z}$ and well determined spectral properties (i.e. \\ep, \\eiso\\ and \\liso) the slope and normalization of these correlations might change.  For comparison we also considered four GRBs with a peak in the GeV light curve. If the GeV emission is interpreted as afterglow (Barniol--Duran \\& Kumar 2009; Ghirlanda et al. 2010; Ghisellini et al. 2010) the measure of $t_{\\rm p,z}$ at early times in the GeV range allows us to estimate their \\G, that are consistent with the correlations found using only the bursts with $t_{\\rm p,z}$ observed in the optical. Although not a proof, this is a hint in favour of the afterglow origin of the GeV emission. % The comoving properties of these bursts are consistent with those of the 27 GRBs with $t_{\\rm p,z}$ measured in the optical % and their inclusion in the sample does not modify substantially the results presented here. However, given the still % debated interpretation of the GeV emission as afterglow we have considered them separately in all our analyses.  These results are schematically summarized in the first column of Tab. \\ref{tab5}. The second column of the same table reports some immediate implications of these results. Since $E^\\prime_{\\rm peak}\\propto E_{\\rm peak}$\\G\\ is contained in a narrow range, all bursts emit their radiation at a characteristic frequency in their comoving frame, irrespective of their bulk Lorentz factor. Furthermore, we can assume that $E_{\\rm peak}\\propto$\\G, and this, together with the quadratic dependence on \\G\\ of $E_{\\rm iso}$ and $L_{\\rm iso}$, yields the ``Amati\" and the ``Yonetoku\" relations. {\\it They are the result of a different \\G--factors.} Indeed, at the extremes of the \\ama\\ and \\yone\\ correlations we find GRB 060218 which has the lowest \\G$\\sim 5$ (inferred from its X--ray and optical properties -- Ghisellini, Ghirlanda \\& Tavecchio  2007), while at the upper end (corresponding to the largest peak energies and isotropic energetics and luminosities) there is GRB 080916C which has the largest \\G=880. The fact that the \\ama\\ and \\yone\\ correlations could be a sequence of \\G\\ factors has been also proposed by Dado, Dar \\& De Rujula (2007) based on different assumptions. %The other intriguing result is the very narrow distribution of $L^\\prime_{\\rm iso}$. %Before discussing if it can be considered as a fundamental quantity or rather a coincidence of %other underlying relations, we can in any case derive important information %about the presence of black body emission. %To this aim, let assume that a sizeable fraction of $L_{\\rm iso}$ is %black body radiation, of comoving temperature $T^\\prime = 10^7 (E^\\prime_{\\rm peak}/4\\, {\\rm keV})$. %In this case: %% %\\begin{eqnarray} %L^\\prime_{\\rm iso} &=& 4 \\pi r^2 \\sigma T^{\\prime 4} \\sim 10^{48} L^\\prime_{\\rm iso, 48} %\\nonumber \\\\ %&\\to& \\, %r = 3.7\\times 10^{11} {L^{\\prime 1/2}_{\\rm iso, 48}\\over T^{\\prime 2}_7} \\,\\, {\\rm cm} %\\end{eqnarray} %% %Since we are dealing with an isotropic equivalent luminosity, %this is the distance at which the black body photons are generated, %and/or the distance at which the fireball becomes transparent. %The narrow distributions of $L^\\prime_{\\rm iso}$ and $E^\\prime_{\\rm peak}$ implies that, %if the black body contributes significantly, then %{\\it $r$ must be nearly the same for all bursts, and that the transparency radius must be smaller than, %or equal to, the above value of $r$}. %This may be difficult to achieve, especially for bursts with the lowest values of \\G\\ %(see Daigne \\& Mochkovitch 2002).   If all bursts had the same jet opening angle, then $L^\\prime_\\gamma =\\theta_{\\rm j}^2 L^\\prime_{\\rm iso}$, and the (logarithmic) width of the $L^\\prime_{\\rm iso}$ distribution would be the same of the (more fundamental) $L^\\prime_\\gamma$ distribution. On the other hand, we have some hints that very energetic and luminous GRBs tend to have narrower opening angles (e.g. Firmani et al. 2005). It is this property that makes the collimation corrected $E_\\gamma$ and $L_\\gamma$ quantities to correlate with $E_{\\rm peak}$ in a different way (i.e. different slope) than in the Amati and Yonetoku relation (Ghirlanda et al. 2004; Nava et al. 2006).  We are then led to propose the following ansatz: the opening angle of the jet inversely correlates with the bulk Lorentz factor $\\theta_{\\rm j}\\propto$ \\G$^{-a}$. There are too few GRBs in our sample with measured $\\theta_{\\rm j}$ to find a reasonable value for the exponent $a$, but it is nevertheless instructive to explore the case $a=1/2$, leading to $\\theta^2_{\\rm j}$\\G$=$ constant. If we assume this relation we find, for the collimation corrected $E_{\\gamma}$: % \\begin{equation} E_{\\gamma} = \\theta^2_{\\rm j} E_{\\rm iso} \\propto \\Gamma_{0} \\propto E_{\\rm peak} \\end{equation} % This is the ``Ghirlanda\" relation in the wind case (Nava et al. 2006). Similarly, for the collimation corrected luminosity (Ghirlanda, Ghisellini \\& Firmani 2006): % \\begin{equation} L_{\\gamma} = \\theta^2_{\\rm j} L_{\\rm iso} \\propto \\Gamma_{0} \\propto E_{\\rm peak} \\end{equation} % Another important consequence of our ansatz is that, in the comoving frame, the collimation corrected energetic $E^\\prime_\\gamma$ becomes constant: % \\begin{equation} E^\\prime_{\\gamma} = \\theta^2_{\\rm j} {E_{\\rm iso}\\over \\Gamma_{0}}=\\,\\,{\\rm constant} \\end{equation} % This allows to ``re--intepret\" the constancy of $L^\\prime_{\\rm iso}$ as a consequence of the constant $E^\\prime_\\gamma$: % \\begin{equation} L^\\prime_{\\rm iso} \\sim {E^\\prime_\\gamma \\over T^\\prime_{90} \\theta^2_{\\rm j}} = {E^\\prime_\\gamma \\over T_{90}  \\theta^2_{\\rm j}\\Gamma_0} = \\,\\,{\\rm constant} \\end{equation} % In other words, in the comoving frame, the burst emits {\\it the same amount of energy at the same peak frequency}, irrespective of the bulk Lorentz factor. For larger \\G\\ the emitting time in the comoving frame is longer (by a factor \\G\\ if the observed $T_{90}$ is the same), so the comoving luminosity is smaller. But since the jet opening angle is also smaller (for larger \\G), the isotropic equivalent luminosity turns out to be the same. These consequences are listed in the third column of Tab. \\ref{tab5}.   Interestingly, we note that the general formula for the estimate of the jet opening angle % \\begin{equation} \\theta_{\\rm j}\\propto \\left(\\frac{t_{\\rm j,obs}}{1+z}\\right)^{{3-s}\\over{8-2s}} \\left(\\frac{n_{0}\\eta}{E_{\\rm iso}}\\right)^{{1}\\over{8-2s}} \\end{equation} % with $s=0$ for the homogeneous case and $s=2$ for the wind case, can be combined with Eq. \\ref{eqgamma} to give: % \\begin{equation} \\theta_{\\rm j}\\Gamma_{0}\\propto \\left(\\frac{t_{\\rm j,obs}}{t_{\\rm p,obs}}\\right)^{{3-s}\\over{8-2s}} \\end{equation} % The product $\\theta{\\rm j}\\Gamma_{0}$ then depends only on two observables, i.e. the time of the peak of the afterglow $t_{\\rm p,obs}$ and the time of the jet break $t_{\\rm j,obs}$, and it is independent from the redshift $z$ and the energetic \\eiso\\ as well as from the density profile normalization $n_{0}$ and radiative efficiency $\\eta$. If also the product $\\theta_{\\rm j}^2\\Gamma_{0}=$const, then we can derive both $\\theta_{\\rm j}\\propto (t_{\\rm p,obs}/t_{\\rm j,obs})^{{3-s}\\over{8-2s}}$ and $\\Gamma_{0}\\propto (t_{\\rm j,obs}/t_{\\rm p,obs})^{{3-s}\\over{4-s}}$. If the ansatz $\\theta_{\\rm j}^2\\Gamma_0=$ const will prove to be true, then by simply measuring the peak time and the jet break time of the afterglow light curve we could estimate both $\\theta_{\\rm j}$ and $\\Gamma_{0}$ for any GRB.   %-------------------------------------------------------------- \\begin{table*} \\begin{center} \\begin{tabular}{llll} \\hline\\hline Our results                        &Implications                                                        &If $\\theta_{\\rm j}^2\\Gamma\\sim$const        \\\\ \\hline $E^\\prime_{\\rm peak}\\sim$ const &$E_{\\rm peak} \\propto\\Gamma$                                           &    \\\\ $E_{\\rm iso} \\propto \\Gamma^2$  &$E_{\\rm iso}\\propto E^2_{\\rm peak}$                                    &$E_\\gamma = \\theta_{\\rm j}^2 E_{\\rm iso}\\propto \\Gamma\\propto E_{\\rm peak}$ \\\\ $L_{\\rm iso} \\propto \\Gamma^2$  &$L_{\\rm iso}\\propto E^2_{\\rm peak}$                                    &$L_\\gamma= \\theta_{\\rm j}^2 L_{\\rm iso}\\propto \\Gamma\\propto E_{\\rm peak}$ \\\\ $T_{90}$ not $f(\\Gamma)$                  &$T^\\prime_{90} \\propto \\Gamma$                                             &$E^\\prime_\\gamma \\sim $ const     \\\\ $L^\\prime_{\\rm iso}\\sim$ const            &$E^\\prime_{\\rm iso}/L^\\prime_{\\rm iso}\\propto T^\\prime_{90}\\propto \\Gamma$ &$L^\\prime_\\gamma \\sim E^\\prime_\\gamma/T^\\prime_{90}\\sim 1/\\Gamma$     \\\\ \\hline \\hline \\end{tabular} \\caption{Schematic summary of our results and their implications for the case of a wind density profile. We have assumed that both $E_{\\rm iso}$ and $L_{\\rm iso}$ scale as $\\Gamma^2$, instead of $\\Gamma^{2.2}$. } \\label{tab5} \\end{center} \\end{table*}  %---------------------------------------------------------------- % ---------------------------------------- \\begin{figure} \\hskip -0.4true cm \\psfig{file=theta_g.ps,height=7cm} \\caption{ Jet opening angle as a function of $\\Gamma_0$ for a H (stars) and for a W (squares). Empty symbols show the jet angles estimated by assuming the consistency of our sample with the $E_{\\rm peak}$--$E_\\gamma$ relation. Filled symbols refer to the bursts of our sample for which the jet opening angle has been calculated from the measured jet break time of the optical light curves. The two lines (dashed for the H case and dot--dashed for the W case) show the powerlaw fit of the data points considering $\\theta_{\\rm jet}$ vs $\\Gamma_{0}$ and $\\Gamma_{0}$ vs $\\theta_{\\rm jet}$ . The gray symbols show the three long bursts with a peak in the GeV light curve that, if interpreted as afterglow emission, allows us to estimate \\G. } \\label{fg8} \\end{figure} % ---------------------------------------- \\begin{figure} \\vskip -0.5 cm \\hskip -0.8true cm \\psfig{file=distribuz_theta_g.ps,height=7cm} \\caption{ Distribution of $\\theta_{\\rm j}\\Gamma_{0}$ in the H and W case (blue and purple histograms) estimated by assuming the $E_{\\rm peak}$--$E_\\gamma$ relation in the H (Ghirlanda et al. 2004) or W (Nava et al. 2006) case. The hatched histograms show the few GRBs in our samples for which $\\theta_{\\rm j}$ has been calculated from the measured jet break time in the optical light curve. } \\label{fg9} \\end{figure} % ----------------------------------------   In our sample, only for 4 bursts we can estimate the jet opening angle from the measure of the jet break time of the optical light curve. Their small number does not make possible to directly test the existence of a relation between $\\Gamma_0$ and $\\theta_{\\rm j}$. However, an estimate of the jet opening angle can be possible by assuming that all bursts in our sample are consistent with the ``Ghirlanda\" relation. Fig. \\ref{fg8} shows the estimated $\\theta_{\\rm j}$ as a function of $\\Gamma_0$. Stars (squares) refers to angles derived under the assumption of a H (W). To estimate the jet opening angles we considered the most updated ``Ghirlanda\" correlation, which comprises 29 GRBs with measured jet break time (Ghirlanda et al. 2006). For the homogeneous density profile the relation has the form $\\log E_{\\rm peak}=-32.81+0.70 \\log E_\\gamma$, while in the case of a W the relation becomes $\\log E_{\\rm peak} =-50.08+1.04 \\log E_\\gamma$. Given the large scatter of the data points in Fig.~\\ref{fg8}, we fitted both $\\theta_{\\rm j}$ versus $\\Gamma_{0}$ and $\\Gamma_{0}$ versus $\\theta_{\\rm j}$: we obtain $\\theta_{\\rm j}\\propto \\Gamma_0^{-0.22}$ and $\\Gamma_0 \\propto \\theta_{\\rm j}^{-2.32}$ for the H case (dashed lines in Fig. \\ref{fg8}) and $\\theta_{\\rm j}\\propto \\Gamma_0^{-0.52}$ and $\\Gamma_0 \\propto \\theta_{\\rm j}^{-1.14}$ for the W case (dot--dashed line in Fig. \\ref{fg8}). We conclude that our ansatz $\\theta_{\\rm j} \\propto \\Gamma_0^{-1/2}$ is consistent with, but not proven by, this analysis.  An interesting exercise is to estimate the product $\\theta_{\\rm j}$\\G. % The opening angle $\\theta_{\\rm j}$ and the bulk Lorentz factor \\G\\ are indeed related. From the observational point of view $\\theta_{\\rm j}$\\GA$\\gg$1 at the end of the prompt phase, so that the decrease of \\GA\\ in the afterglow phase, due to the interaction of the GRB fireball with the interstellar medium, gives rise to a jet break when $\\theta_{\\rm j}$\\GA$\\sim$1.  Some numerical simulations (Komissarov et al., 2009) of jet acceleration have shown that a magnetic dominated jet confined by an external medium should have $\\theta_{\\rm j}$\\G$\\le 1$. This value is inconsistent with typical values of $\\theta_{\\rm j}$ and \\G: in the case of an homogeneous wind density profile the typical $\\theta_{\\rm j}\\sim 0.1$ radiants (Ghirlanda et al. 2007) while in the case of a wind density profile $\\theta_{\\rm j}\\sim 0.07$ radiants. Combining these values with the average values of \\G\\ estimated in this paper (Tab. \\ref{tab1}) we find $\\theta_{\\rm j}$\\G$\\sim 14$ (5) for the H (W) case.  These are approximate values: the sample of GRBs with measured $\\theta_{\\rm j}$ (Ghirlanda et al. 2007) contains only 4 bursts of the sample of  events of the present paper with estimated \\G. However, though somehow speculative, we can derive $\\theta_{\\rm j}$ for the 32 GRBs of our sample assuming the \\ghi\\ correlation in the H case (Ghirlanda et al. 2004) or in the W (Nava et al. 2006). In Fig. \\ref{fg9} we show the distributions of the product $\\theta_{\\rm j}$\\G\\ in the H case (blue histogram) and in the W case (purple histogram). We note that both are centered around typical values of 20 and 6 (for the H and W case, respectively). These values are in good agreement with the results of recent simulations of (i) a magnetized jet confined by the stellar material that freely expands when it breaks out the star (Komissarov, Vlahakis \\& Koenigl  2010) or (ii) a magnetized unconfined split--monopole jet (Tchekhovskoy, McKinney \\& Narayan 2009; Tchekhovskoy, Narayan \\& McKinney 2010). A possible test of these two scenarios could be short GRBs where the absence of the progenitor star would prefer model (ii) for the jet acceleration. In our sample only the short/hard GRB 090510 is present. No jet break was observed for this event and in general we do not yet know if short GRBs follow the same \\ghi\\ correlation of long ones."
    },
    "1107/1107.2308_arXiv.txt": {
        "abstract": "Neutrino-neutrino interactions in dense neutrino streams, like those emitted by a core-collapse supernova, can lead to self-induced neutrino flavor conversions. While this is a nonlinear phenomenon, the onset of these conversions can be examined through a standard stability analysis of the linearized equations of motion. The problem is reduced to a linear eigenvalue equation that involves the neutrino density, energy spectrum, angular distribution, and matter density. In the single-angle case, we reproduce previous results and use them to identify two generic instabilities: The system is stable above a cutoff density (``cutoff mode''), or can approach an asymptotic instability for increasing density (``saturation mode''). We analyze multi-angle effects on these generic types of instabilities and find that even the saturation mode is suppressed at large densities. For both types of modes, a given multi-angle spectrum typically is unstable when the neutrino and electron densities are comparable, but stable when the neutrino density is much smaller or much larger than the electron density. The role of an instability in the SN context depends on the available growth time and on the range of affected modes. At large matter density, most modes are off-resonance even when the system is unstable. ",
        "introduction": "\\label{sec:intro}  Neutrino flavor oscillations in a supernova (SN) are strongly suppressed by matter effects~\\cite{Wolfenstein:1977ue} until the neutrinos pass through the usual MSW region~\\cite{Mikheev:1986gs, Mikheev:1986if, Dighe:1999bi, Kuo:1989qe} far out in the envelope of the collapsing star. However, neutrino-neutrino interactions~\\cite{Pantaleone:1992eq, Sigl:1992fn}, through a flavor off-diagonal refractive index, can trigger self-induced flavor conversions~\\cite{Samuel:1993uw, Kostelecky:1993dm, Kostelecky:1995dt, Samuel:1996ri, Sawyer:2005jk, Sawyer:2008zs}. This collective effect tends to occur between the neutrino sphere and the MSW region and can lead to strongly modified neutrino spectra, showing features such as spectral swaps and splits~\\cite{Duan:2006an, Raffelt:2007cb, Duan:2007fw, Fogli:2007bk, Fogli:2008pt, Dasgupta:2009mg}; for a review see Ref.~\\cite{Duan:2010bg}. The overall scenario, supported by heuristic arguments and numerical examples, is that deep inside the SN core, the system performs ``synchronized oscillations'' with an extremely small amplitude, i.e.\\ every neutrino remains essentially stuck in its initial flavor eigenstate. As the neutrinos stream outwards, there is a sharp onset radius where ``bimodal'' oscillations begin: Some ranges of modes start pendulum-like oscillations~\\cite{Samuel:1996ri, Hannestad:2006nj, Duan:2007mv, Raffelt:2011yb}, exchanging their flavor content with each other without affecting the flavor content of the overall system.  This scenario engenders a crucial simplification for the treatment of neutrino transport in SN simulations. At high densities, where neutrinos collide frequently, it is enough to solve the transport equations for each flavor separately, ignoring oscillations entirely. On the other hand, flavor conversions at larger distances can be treated ignoring neutrino collisions and absorption, i.e.\\ as a pure propagation problem. So collisions and flavor oscillations are phenomena that are assumed to be taking place in different regions of the star and can be treated independently. When the radial distance where bimodal oscillations begin is far away from the SN core, this assumption is valid and the flavor conversions do not affect the SN dynamics. Recent studies dedicated to the SN accretion phase, under simplifying assumptions, once more confirm this picture~\\cite{Chakraborty:2011gd, Dasgupta:2011jf}.  However, what is missing is a systematic approach to decide, without solving the equations of motion, if self-induced flavor conversions occur for given neutrino spectra (flavor-dependent energy and angular distribution), overall neutrino density, and matter density. Formal stability criteria exist only in the ``single-angle approximation'' where it is assumed that all neutrinos feel the same neutrino-neutrino refractive effect. In this case the analytic pendulum solution has been found and its existence and parameters can be calculated from the neutrino spectrum and density alone~\\cite{Dasgupta:2009mg}.  On the other hand, the current-current nature of the low-energy weak-interaction Hamiltonian implies that neutrinos in the background of an anisotropic neutrino flux experience a refractive effect that strongly depends on direction. For some energy spectra, these ``multi-angle effects'' have little impact, whereas in other cases they completely change the solution. A SN neutrino flux with a vanishing net $\\nu_e$ flux is unstable everywhere and shows quick multi-angle decoherence~\\cite{Raffelt:2007yz}. If the $\\nu_e$ flux is large enough, this effect is avoided if the energy spectrum is simple~\\cite{EstebanPretel:2007ec}. More complicated spectra can be unstable even for a large $\\nu_e$ flux at any density in the single-angle case, but the instability is suppressed by multi-angle effects~\\cite{Raffelt:2008hr, Duan:2010bf}. In addition, the presence of ordinary matter causes a multi-angle suppression of the bimodal instability~\\cite{EstebanPretel:2008ni}. On the other hand, it was claimed that for nontrivial angular distributions, as may exist when different flavors are emitted from very different neutrino spheres, there is a novel multi-angle instability~\\cite{Sawyer:2005jk}. Deviations from the usually assumed cylindrical symmetry may also have an important influence~\\cite{Sawyer:2008zs}.  Although our problem is nonlinear and therefore would seem intractable, noting that an instability must occur in order for the onset to take place leads to a surprising simplification. In the dense SN matter well inside of the MSW region, the matter effect is so large that neutrino propagation eigenstates are essentially identical with flavor eigenstates. This means that in the weak-interaction basis, the flavor matrices of occupation numbers are almost perfectly diagonal, allowing us to linearize the equations of motion (EoMs) in terms of the small off-diagonal elements. An instability is equivalent to some of these small elements starting to grow exponentially. All we need to do is linearize the EoMs, perform a Fourier analysis, and seek exponentially growing solutions of the relevant eigenvalue equation. This is simple even in the multi-angle situation. Almost certainly this approach can be extended to cases without the usually assumed cylindrical symmetry around the radial direction.  Studying the stability of a strongly coupled and nonlinear system in the small-amplitude limit is a standard technique. In the present context it was put forth in Ref.~\\cite{Sawyer:2008zs}. However, the method was carried only to schematic cases of a small number of neutrino momentum modes, leaving open how to apply it to realistic situations.  While a stability analysis provides crucial insight, we stress that it alone is not enough to assess the impact in the SN context. Only a small range of modes may participate in the bimodal oscillation or the growth rate may be too small on the available time scale. (Of course, both the growth rate and the location of the resonance on the energy-angle spectrum are found by solving the eigenvalue equation.) For example, the classic single-angle case with an initial Fermi-Dirac spectrum of only $\\nu_e$ and $\\bar\\nu_e$ is always unstable and the usual concept of a synchronization radius does not apply. However, for large neutrino densities, the growth rate is small and only a narrow range of infrared modes is affected. A ``visible'' effect arises only at a quasi-onset radius where the growth rate begins to compete with the overall evolution time scale and the resonance begins to move into the main part of the spectrum. We will here largely avoid such issues and concentrate on setting up the method and discussing simple but informative schematic cases. Applying this method in a realistic SN context will be left to future work.  Our study in Sec.~\\ref{sec:eoms} leads to the linearized equations of motion at a large distance from the neutrino source in the two-flavor case, with an azimuthally symmetric neutrino emission. In Sec.~\\ref{sec:single}, we present the stability analysis in the single-angle approximation, and illustrate it in Sec.~\\ref{sec:single-ex} with the examples of box spectra where the results may be understood analytically. In Sec.~\\ref{sec:nosleep} we point out some special features of realistic spectra that do not vanish at low energies. Sections~\\ref{sec:multi} and \\ref{sec:multi-ex} demonstrate how the single-angle results are modified by the inclusion of multi-angle and matter effects. The latter one also analyzes a realistic SN spectrum using the insights obtained from the box spectra. In Sec.~\\ref{sec:novel}, we show that a multi-angle spectrum with a zero crossing is unstable in both hierarchies if the lepton asymmetry is small. In Sec.~\\ref{sec:conclusions} we conclude with a brief summary of our findings and an outlook on future directions. ",
        "conclusions": "\\label{sec:conclusions}  The refractive effect of neutrino-neutrino interactions can cause a flavor instability among the different energy and angular modes, exchanging the flavor content while leaving unchanged the overall flavor content of the entire ensemble. In the context of SN neutrinos, the focus so far has been on impacts of collective flavor oscillations on the neutrino signal of the next nearby SN or for the r-process nucleosynthesis. Many analyses have been performed under the simplifying assumption that all neutrinos feel the same refractive effect due to the other neutrinos (single-angle assumption). However it has also been found that multi-angle effects can strongly modify the answer and can typically suppress flavor conversions that would occur in the single-angle treatment. In this sense, a more pressing question is to understand the stability conditions for a dense neutrino stream, and to understand if the usual neglect of flavor oscillations in those dense SN regions is justified where a Boltzmann treatment of neutrino transport is necessary.  We have studied the conditions under which instabilities can occur in the neutrino stream in the two-flavor framework. While the equations of motion are nonlinear, the run-away solutions can be found by a standard mode analysis of the linearized system, i.e.\\ in the small-amplitude expansion. We have shown how the stability question is reduced to a straightforward eigenvalue equation in the co-rotating frame. The existence of complex eigenvalues directly corresponds to an instability. In the multi-angle case, besides the small-amplitude expansion, we have linearized the equations in the small angular divergence of the neutrino flux that is applicable at a large distance from the source. In the single-angle case, our conditions agree with the previous literature, whereas in the multi-angle case, our results are new.  In the single-angle case and using schematic energy spectra (box spectra in the $\\omega$ variable), the eigenvalue equation can be solved analytically. We identify two generic instability cases. The ``cutoff mode'' exists for a finite range of neutrino densities, or equivalently, a finite interval of effective neutrino-neutrino interaction strengths $\\mu_0<\\mu<\\mu_1$. In particular, the system is stable above a cutoff value $\\mu_1$, also termed the ``synchronization'' or ``sleeping top'' regime. This behavior is generic to a single-crossed (or two-box) spectrum and represents the traditional flavor pendulum. On the other hand, there is the ``saturation mode'' where the instability approaches an asymptotic value for large $\\mu$. This mode requires at least two crossings (three boxes) and had been previously found in the literature. When it exists, the saturation mode exists for both hierarchies. For multi-crossed spectra, several modes can co-exist, for example in the four-box case a saturation mode and a cutoff mode or two cutoff modes.  An additional case occurs for neutrino spectra with a vanishing asymmetry, i.e.\\ with no net lepton number flux. The run-away rate $\\kappa$ grows indefinitely as a function of $\\mu$. This actually is a special case of the cutoff mode (cutoff at infinite $\\mu$) and does not seem to be of practical interest in the SN context.  Realistic neutrino spectra are not described by boxes and have tails as a function of $\\omega=|\\Delta m^2|/2E$, corresponding to infrared energies. Such spectra do not have a cutoff and a synchronization regime does not exist. However, for large $\\mu$ the instability moves entirely into the infrared regime and has no practical impact. There remains an effective cutoff density where the amount of flavor exchange between modes drops by orders of magnitude within a narrow range of neutrino densities. Still, the existence of a synchronization regime is only an approximate concept.  Multi-angle effects modify the instability in several ways. For example, while the single-angle approximation predicts an instability only if there is a positive zero-crossing of the spectrum $g_\\omega$, the multi-angle analysis implies that spectra with a small asymmetry (a small net lepton number flux) and in the absence of matter are unstable in both hierarchies, whatever be the sign of the zero-crossing. The solutions for both hierarchies can be mapped onto each other.  The main multi-angle modification of the instability is caused by a background of matter, where here neutrinos themselves also contribute to the matter effect. The saturation mode is ``fragile'' and, for box spectra, develops a cutoff. (For realistic spectra with tails, there is no exact cutoff, however, as mentioned earlier.) The neutrino background alone is enough to achieve this effect.  For both the saturation and cutoff modes, a large matter effect shifts the domain of instability $\\mu_0<\\mu<\\mu_1$ to larger values of $\\mu$. The amount of shift is approximately equal to the matter potential $\\lambda$, apparently without reducing the typical $\\kappa$ values within the instability domain. An additional consequence is that the affected range of modes shrinks for increasing $\\lambda$. Even if the run-away scale $\\kappa$ is not reduced, only a narrow range of angular modes is on resonance if $\\lambda$ is large. For realistic spectra with tails, $\\kappa(\\mu)$ does not vanish for large $\\mu$, i.e.\\ there is not an exact cutoff. On the other hand, there is a lower threshold $\\mu_0\\sim\\lambda$ below which the spectrum is completely stable. In other words, multi-angle effects introduce a range $0<\\mu<\\mu_0\\sim\\lambda$ where the spectrum is stable even when such a stable range does not exist in the single-angle approximation. This behavior corresponds to the multi-angle suppression of instability at $\\lambda \\gg \\mu$, which we motivate analytically through our linearized treatment.  One may further investigate these findings in the context of realistic SN situations where neutrinos stream from regions of large density to vacuum, sweeping through $\\mu$ values from very large to zero. The central frequency $\\gamma$ of the instability then sweeps through different spectral ranges and the growth rate $\\kappa$ depends on both the neutrino and matter density. Effective flavor conversion can be suppressed in different ways: The growth rate $\\kappa$ can vanish when $\\mu$ falls into the multi-angle-implied stability domain, or it can be too small on the available distance scale, or the instability can be located mostly in the infrared and affect only an ignorable part of the spectrum.  We have considered neutrinos streaming from a source and have made the approximation of small angular divergence that is appropriate at a large distance from the source. Our approach can be easily adapted to a homogeneous ensemble with a nontrivial angular distribution that evolves in time rather than as a function of distance from a source. In this case no expansion in the angular variable is needed.  We have only considered trivial angular distributions where all neutrino flavors are assumed to be emitted from a common blackbody sphere without limb darkening, corresponding to a box spectrum in our $u$ variable. One can extend our analysis to more complicated angular distributions that could involve, for example, spectral crossings in the angular variable. As advocated by Sawyer, such crossings cause a ``multi-angle instability,'' but its practical relevance remains to be investigated.  We have made the usual assumption of azimuthal symmetry of the neutrino flux around the radial direction away from the source. In the long-distance limit of small angular divergence of the neutrino field, one can of course expand the EoMs in terms of small variables that depend on two angles. It remains to be explored if cylindrical symmetry enforces unphysical solutions, just as the single-angle approximation sometimes enforces unphysical solutions, and if there are new instabilities or important modifications that could be relevant in the SN context. Based on schematic models using a small number of modes it has been suggested that violations of cylindrical symmetry could have a major impact~\\cite{Sawyer:2008zs}.  A rich phenomenology of neutrino flavor conversions is bound to emerge from further investigations into the collective effects due to neutrino-neutrino interactions and ordinary matter background, in the complete multi-angle framework."
    },
    "1107/1107.4582_arXiv.txt": {
        "abstract": "In this paper we derive, directly from the Nambu-Goto action, the relevant components of the acceleration of cosmological featureless $p$-branes, extending previous analysis based on the field theory equations in the thin-brane limit. The component of the acceleration parallel to the velocity is at the core of the velocity-dependent one-scale model for the evolution of $p$-brane networks. We use this model to show that, in a decelerating expanding universe in which the $p$-branes are relevant cosmologically, interactions cannot lead to frustration, except for fine-tuned non-relativistic networks with a dimensionless curvature parameter $k \\ll 1$. We discuss the implications of our findings for the cosmological evolution of $p$-brane networks. ",
        "introduction": "It is generally accepted that the universe underwent a phase of accelerated expansion in its early history. This paradigm, usually  denominated cosmological inflation \\cite{Guth:1980zm,Linde:1981mu}, explains the extreme flatness and homogeneity of the observed universe, as well as the the origin of the large scale structure. In the context of the brane-world realization of string theory, cosmological inflation could be driven by the interaction between p-dimensional D-branes \\cite{Dvali:1998pa,Burgess:2001fx,ALexander:2001ks} (which are, along with fundamental strings, the fundamental objects of the theory). Brane inflationary  scenarios typically end with a symmetry breaking phase transition, leading to the production of branes with lower dimensionality. The copious production of $1$-dimensional branes (cosmic strings) is predicted for a large variety of brane inflation models, while higher-dimensional $p$-branes might also be produced \\cite{Sarangi:2002yt,Majumdar:2002hy,Jones:2002cv}. Therefore, brane inflation may lead to the production of $p$-branes networks, evolving in higher dimensional backgrounds.  Cosmic strings and domain walls are the only non-trivial $p$-brane solutions allowed in $3+1$ dimensional backgrounds. The interest in cosmic strings has recently been revived due to the possibility of detecting their cosmological signatures observationally. In particular, they may leave an observable imprint on the B-mode polarization \\cite{Seljak:2006hi,Pogosian:2007gi} and small-scale anisotropy \\cite{Pogosian:2008am} of the  cosmic microwave background which may be within reach of the Planck mission. Moreover, there is also the prospect of the future detection of the gravitational waves emitted by cosmic strings with the LIGO2 and LISA probes \\cite{Polchinski:2007rg} (see also \\cite{Avelino:2003nn,Morganson:2009yk,Brandenberger:2010hn} for other examples of possible observational signatures of cosmic strings). In order to make precise observational predictions it is necessary to understand the nature of cosmic strings and their late-time evolution.  While standard cosmic string networks have been extensively studied both using a semi-analytical velocity-dependent one-scale (VOS) model \\cite{VOS} and numerical simulations \\cite{Allen:1990tv}, significantly less is known about the dynamics of complex string networks with junctions. For example, cosmic superstring networks may have an hierarchy of tensions and junctions \\cite{Copeland:2003bj} and, consequently, their observational signatures are expected to be different from those of ordinary cosmic strings \\cite{Brandenberger:2008ni,Danos:2009vv}. Still, it is not yet clear whether or not these differences prevent cosmic superstring networks from attaining a linear scaling regime and, therefore, the late-time evolution of these networks is not yet completely understood \\cite{Copeland:2005cy,Avgoustidis:2007aa,Pourtsidou:2010gu}.  A $p$-brane network, if frozen in comoving coordinates (or frustrated), is characterized by a negative average pressure. This property is particularly interesting in the case of domain wall networks. There is compelling evidence \\cite{Komatsu:2010fb,Amanullah2010} which indicates that the universe is currently undergoing a phase of accelerated expansion, caused by an exotic energy component (dubbed dark energy) accounting for more than two thirds of the energy density of the universe and whose nature is yet unknown. If domain walls were the dominant energy component, and provided their root-mean-square (RMS) velocity was small enough, they could drive a phase of accelerated expansion. For this reason, frustrated domain wall networks have been suggested as a dark energy candidate \\cite{Bucher:1998mh}. The conditions under which domain wall networks may frustrate as a result of cosmological evolution were intensively studied in \\cite{PinaAvelino:2006ia,Avelino:2006xy,Battye:2006pf,Avelino:2006xf,Avelino:2008ve,Sousa:2009is} using a semi-analytical VOS model and field-theory simulations. These studies indicate that frustration does not result naturally from the evolution of realistic and cosmologically relevant domain wall networks.  The cosmological dynamics of generic $p$-brane networks has been studied in \\cite{Sousa:2011ew} using a VOS model characterizing the evolution of their RMS velocity and characteristic length. As in the case of cosmic string \\cite{Martins:1996jp} and domain wall networks \\cite{Avelino:2006xf,Avelino:2011ev}, in homogeneous and isotropic universes with a decelerating power law expansion, $p$-brane networks may evolve towards a linear scaling solution in which the characteristic length grows proportionally to the Hubble radius and the RMS velocity remains constant \\cite{Sousa:2011ew}. This scale-invariant solution is an attractor of the VOS equations, characterizing the late-time evolution of the networks in a frictionless regime.  In this paper we derive the evolution equation for the velocity of featureless infinitely-thin $p$-brane networks in homogeneous and isotropic universes with an arbitrary number of spatial dimensions, directly from the Nambu-Goto action. This generalizes previous results obtained using field theory equations in the thin-brane limit, further validating the VOS model for the evolution of $p$-brane networks in $N+1$-dimensional homogeneous and isotropic universes derived in \\cite{Sousa:2011ew}. We also consider the effect of a generic interaction mechanism between the $p$-branes and other cosmological components, studying its  potential role in the frustration of the networks.  Throughout the work, we will assume the metric signature $[+,-,...,-]$ and the calculations will be done using units in which $c=1$. The Einstein summation convention will be used when a greek index variable appears twice in a single term, once in an upper (superscript) and once in a lower (subscript) position. ",
        "conclusions": "}  In this paper, we derived the equation of motion for infinitely thin featureless $p$-branes, by computing the tangential and normal components of the acceleration directly from the Nambu-Goto action. Our results further validate the semi-analytical VOS model developed in \\cite{Sousa:2011ew}, and its use in dynamical studies of $p$-brane networks. The VOS model unifies, in a single framework, the dynamics of $p$-brane networks for any possible values of the pair $(N,p)$. While part of the dynamical dependency on the parameters $N$ and $p$ is explicit in the equations of motion, there is also an implicit dependency in the parameters ${\\tilde c}$, $\\ell_f$ and $k$. In this paper we demonstrated that, if the $p$-branes are the dominant component of the universe, then frustration is not possible except if the curvature parameter is driven towards very small values for non-relativistic networks or if the expansion is accelerated. In the case of domain walls there is very strong analytical and numerical evidence (both in two $(N=2,p=1)$ and three $(N=3,p=2)$ spatial dimensions) that $k$ is never becomes much smaller than unity (except deep into inflationary or friction dominated regimes), thus preventing frustration from being attained, at least for realistic domain wall networks playing a dark energy role.  We conjecture that this may be a general result, valid for any realistic $p$-brane network independently of the values of $N$ and $p$ with $1 \\le p \\le N-1$."
    },
    "1107/1107.1969_arXiv.txt": {
        "abstract": "{ } {We report spectroscopic observations in support of a novel view of transition region explosive events, observations that lend empirical evidence that at least in some cases explosive events may be nothing else than spinning narrow spicule-like structures.} { Our spectra of textbook explosive events with simultaneous Doppler flow of a red and of a blue component are extreme cases of high spectroscopic velocities that lack apparent motion, to be expected if interpreted as a pair of collimated, linearly moving jets. The awareness of this conflict led us to the alternate interpretation of redshift and blueshift as spinning motion of a small plasma volume. In contrast to the bidirectional jet scenario, a small volume of spinning plasma would be fully compatible with the observation of flows without detectable apparent motion. We suspect that these small volumes could be spicule-like structures and try to find evidence. We show observations of helical motion in macrospicules and argue that these features -- if scaled down to a radius comparable to the slit size of a spectrometer --  should have a spectroscopic signature similar to that observed in explosive events, while not easily detectable by imagers. Despite of this difficulty, evidence of helicity in spicules has been reported in the literature. This inspired us to the new insight that the same narrow spinning structures may be the drivers in both cases, structures that imagers observe as spicules and that in spectrometers cross the slit and are seen as explosive events. } { We arrive at a concept that supports the idea that explosive events and spicules are different manifestations of the same helicity driven scenario. In contrast to the conventional view of explosive events as linear bidirectional jets, that are triggerred by a reconnection event in the transition region, this new interpretation is compatible with the observational results. Consequently, in such a case, a photospheric or subphotosperic trigger has to be assumed. } {We suggest that explosive events / spicules are to be compared to the unwinding of a loaded torsional spring.} ",
        "introduction": "Explosive events (EEs) are characterized as short-lived, small-scaled incidents of rapid plasma acceleration to typically 50~km/s\\, to \\,150~km/s and sometimes even higher velocities in both directions. In spectroscopic data EEs are easily detected by the redshift and blueshift of the observed transition region (TR) line. The terminology 'explosive event' has first been introduced by \\citet{Dere84} based on the analysis of high-resolution spectra of TR emission lines obtained by the HRTS instrument on {\\it Spacelab}, but it turned out that this term is quite debatable \\citep[e.g.,][]{Dere89} and may be misleading.% With the advent of {\\it SoHO}-SUMER \\citep{Wilhelm95} a revival of this field of research started. Typical EEs were found to be short-lived (60~s to 200~s), small-scale (1500~km to 2500~km), high-velocity ($\\pm$50~km/s to $\\pm$150~km/s) flows that occur very frequently, sometimes in bursts. \\citet{Teriaca04} estimated an average size of 1800~km, a birth rate of 2500~s$^{-1}$, and 30,000 events at any one time on the entire Sun. In the classical view EEs are seen as bi-directional jets that are generated by a Petschek-type reconnection event \\citep{Innes97} high up in the TR with a collimated upflowing component -- blueshifted in TR emission -- and with a downflowing, redshifted component at some angle to the line-of-sight (LOS). Statistically, the blueshift dominates the redshift in magnitude \\citep[e.g.,][]{Innes97,Madjarska02,Zhang10}, a fact that is explained by the deceleration of the downflowing material hitting denser atmospheric layers. \\citet{Ning04} found that EEs tend to cluster near regions of evolving network fields and speculated that the periodicity of 3\\,-\\,5 minutes found in EE bursts may be related to subsurface phenomena. An overview of the immense amount of work on small-scale dynamical events and related cross-references are found in the review of \\citet{Innes04}.  From the very beginning, however, there has been the problem that apparent motion as the result of such high-velocity events has not been observed by imaging instruments. \\citet{Dere89} already noted {\\it '... the lack of detectable apparent motions of such high-velocity events'}. \\citet{Innes04} again mention this conflict that is still unsolved. It is the central point of our work to spur a discussion on this discrepancy.  It is intriguing to see that other solar phenomena exist with very similar characteristics in terms of size, duration, temperature, occurrence rate, light curves, or repeatability. \\citet{Madjarska03} tried to establish a relationship with limb phenomena, that are observed with a different geometry, and suggested that blinkers (as observed by CDS) are the on-disk signature of spicules. While in this article they still assume that blinkers and EEs are two separate phenomena not directly related or triggering each other, they later \\citep{Madjarska09} state that {\\it 'the division of small-scale transient events into a number of different subgroups, for instance EEs, blinkers, spicules, surges or just brightenings, is ambiguous'}. Also \\citet{Wilhelm00} argues that there seems to be no obvious distinction between macrospicules and other spicules, apart from the fact that macrospicules are restricted to coronal hole locations. With this proposition, \\citet{Madjarska06} conclude that blinkers, EEs and macrospicules are indeed identical phenomena that are observed with different instruments and with a different geometry. They, however, still assume flows of rising and falling plasma.  \\citet{Innes99} present a 2D-reconnection model for EEs, yet they already mention the possibility of an alternate interpretation of spinning plasma. This ambiguity of a configuration as bi-directional jet or as a spinning volume of gas is explicitely mentioned by \\citet{Innes01}, but still unsolved. The dynamical events presented in the following sections clearly favour the latter explanation.  Helicity is often observed in large-scale events like coronal mass ejections, and is also found in small-scale phenomena like coronal bright points and X-ray jets \\citep[e.g.,][]{Shen11,Liu11}. \\citet{Tian08} found that the Doppler shift pattern of a coronal bright point (BP) gradually varies with height, suggesting that the magnetic loops associated with the BP are twisted or in helical form. Recently \\citet{Kamio10} communicated the observation of an X-ray jet with helical motion at TR temperatures observed by both spectrometers {\\it Hinode}-EIS and {\\it SoHO}-SUMER. \\citet{Nistico09} report in their survey of {\\it STEREO}-EUVI jets that 31 out of 79 events exhibit helical motion, and further mention the possibility that the rest were very narrow so that possibly the twist could not be resolved. This notion has recently been supported by \\citet{Sterling10}, who suggest that macroscopic coronal jets can be scaled down to spicule-size features.  We present a case study of two EEs to demonstrate the conflict with the standard bi-directional jet model and to stress the enigmatic discrepancy of lacking apparent motion. As a solution of this conflict we suggest -- backed up by observational evidence for helicity in spicules -- that a spinning motion may be the source of EEs. Combining the hypothetical concept of spinning spicules with the observation of quasi-stationary Doppler-flow in EEs could be the solution of the conflict. ",
        "conclusions": "We propose a hypothesis suggesting that disk EEs and limb spicules are the same phenomenon. However, we assume an alternateive configuration to explain the flows. It is clear that helicity is behind both manifestations and should be included to understand the physical nature of EEs. The assumption of a swirling narrow cylindrical body, rotating while upflowing, can reproduce observations by both imagers and spectroscopes which can explain the discrepancy between spectroscopic motion and apparent motion. Also, the statistical blueshift dominance of EEs would be an obvious consequence in such a scenario  (see Fig.4, case c). Although our observations do not strictly exclude the possibility of bi-directional jets, there are good reasons to assume that EEs are indeed the spectroscopic signature of spinning type II spicules crossing the spectrometer slit in many cases. Even more, there seems to be neither an obvious distinction between macrospicules and microspicules nor between blinkers and EEs.  The swirling component is normally not detectable by filtergraph instruments, but adds considerably to the energy released by the apparent motion that is detected by such instruments. We note that the scenario of spinning spicule-like features that are rooted in the photosphere requires photospheric or even subsurface sources which is not compatible with the model of reconnection events in the TR. This aspect may require the distinction of different types of EEs and calls for more systematic work. We speculate that the helicity may be related to global helicity as generated by the differential rotation. Alternatively, local reconnection due to the subsurface turbulence in twisted flux tubes as discussed in MHD models could be a possible driving mechanism. In this context, it would be worthwhile to study the chirality of the events and look for different preferences in both hemispheres. These hypotheses are, however, not supported by our data and beyond the scope of this communication.  The IRIS mission -- providing fast spectroscopic capabilities complemented by a chromospheric imager -- will be an ideal platform for systematic statistical analyses of geometrical effects and their imprints on the center-to-limb variation of red-blue tilt, red-blue asymmetry, birth rate, and mean velocity of EEs and also assessing the dominance of helicity in EEs in a quantitative manner."
    },
    "1107/1107.1191_arXiv.txt": {
        "abstract": "{Dispersive Alfven waves (DAWs)  offer, an alternative to magnetic reconnection, opportunity to accelerate solar flare particles in order to alleviate the problem of delivering flare energy to denser parts of the solar atmosphere to match X-ray observations.} {Here we focus on the effect of DAW polarisation, left, right, circular and elliptical, in the different regimes inertial and kinetic, aiming to study these effects on the efficiency of particle acceleration.} {We use 2.5D particle-in-cell simulations to study how the particles are accelerated when DAW, triggered by a solar flare, propagates in the transversely inhomogeneous plasma that mimics solar coronal loop.} {(i) In the inertial regime, fraction of accelerated electrons (along the magnetic field), in the density gradient regions is $\\approx 20\\%$ by the time when DAW develops three wavelengths and  is increasing to $\\approx 30\\%$ by the time when DAW develops thirteen wavelengths. In all considered cases ions are heated in the transverse to the magnetic field direction and fraction of heated particles is $\\approx 35 \\%$. (ii) The case of right circular, left and right elliptical polarisation DAWs, with the electric field in  the non-ignorable transverse direction exceeding several times that of in the ignorable direction, produce more pronounced parallel electron beams (with larger maximal electron velocities) and transverse ion beams in the ignorable direction. In the inertial regime such polarisations yield the fraction of accelerated electrons  $20\\%$. In the kinetic regime this increases to $35\\%$. (iii) The parallel electric field that is generated in the density inhomogeneity regions is independent of the electron-ion mass ratio and stays of the order $0.03\\omega_{pe}c m_e /e$ which for solar flaring plasma parameters exceeds Dreicer electric field by eight orders of magnitude. (iv) Electron beam velocity has the phase velocity of the DAW. Thus electron acceleration is via Landau damping of DAWs. For the Alfven speeds of $V_A=0.3c$ the considered mechanism can accelerate electrons to energies circa 20 keV. (v) The increase of mass ratio from $m_i/m_e=16$ to 73.44 increases fraction of accelerated electrons from $20\\%$ to $30-35\\%$ (depending on DAW polarisation). For the mass ratio $m_i/m_e=1836$ the fraction of accelerated electrons would be $>35\\%$. (vi) DAWs generate significant density and temperature perturbations that are located in the density gradient regions. } {DAWs propagating in the transversely inhomogeneous plasma can effectively accelerate electrons along the magnetic field and heat ions across it.} ",
        "introduction": "Charged particles, accelerated to the speeds exceeding local thermal velocity, play a vital role in numerous laboratory, space, solar and astrophysical interdisciplinary processes. For example, in space plasmas, the Van Allen radiation belts form a torus of energetic charged plasma particles around Earth, which is held in place by Earth's magnetic field. These accelerated, charged particles come from cosmic rays. The radiation belts pose serious hazard to the space technology and future manned space travel. The outer electron radiation belt is mostly produced by the inward radial diffusion \\cite{2004GeoRL..3108805S} and local acceleration \\cite{2005Natur.437..227H} due to transfer of energy from whistler mode plasma waves to radiation belt electrons. Dispersive Alfven waves (DAWs) are known to accelerate electrons in the Earth auroral zone \\cite{2000SSRv...92..423S}. In solar coronal plasmas, solar flare models involve magnetic reconnection which converts the magnetic energy into heat and kinetic energy of accelerated particles. About 50-80\\% of the energy released during solar flares is converted into the energy of accelerated particles \\cite{2004JGRA..10910104E}. Thus, knowledge of the details of particle acceleration via magnetic reconnection (which is deemed a major mechanism operating solar and stellar flares) is of great importance. This is of direct relevance to the solar activity that affects humankind via hazardous phenomena such as solar flares and energetic particles. In astrophysical plasmas, accelerated particles play a major role in explaining the existence of soft excess emission originating from clusters of galaxies that is of fundamental cosmological importance \\cite{2008SSRv..134..207P}. Also, multi-wavelength observations of Active Galactic Nuclei have now provided evidence that they are site of production of high temperature plasmas ($kT > 100$ MeV) with strong tails of accelerated particles. The generation of hot, relativistic plasmas and accelerated particles is a challenging problem both in the context of Active Galactic Nuclei and the origin of high-energy cosmic rays \\cite{1984AdSpR...4..345F}. In laboratory plasmas, in the case of ions, plasma regimes in which enhanced fast-ion transport occurs because of strongly driven energetic particle models (EPMs) should be avoided in order to reach reactor relevant performance and prevent first wall damage \\cite{2006PPCF...48B..15Z}. The investigation of the transition from regimes characterised by weakly driven modes (such as Alfven Eigenmodes, e.g. toroidal Alfven eigenmodes, TAEs) to regimes dominated by strongly driven modes (like EPMs) and the determination of the related threshold are then of crucial importance for burning plasma operations \\cite{2009NucFu..49g5024V}. In the case of electrons, e.g. in the laser driven inertial confinement fusion, accelerated electrons carry away a lot of energy pumped by the laser. Also, due to the laser's strong coupling with hot electrons, premature heating of the dense plasma (ions) is problematic and fusion yields are low. Significant improvements in the modelling of these phenomena have been achieved \\cite{2010PhPl...17e6702K}. Edge Localised Modes (ELMs) are MHD instabilities destabilised by the pressure gradient in the H-mode Tokamak edge pedestal. ELMs produce losses of up to 10\\% plasma energy in several 100 micro-seconds. ELMs are major concern for the operation of ITER. Thus, ELM control is essential. ELMs are the focus of increasing attention by the edge physics community because of the potential impact that the large divertor heat pulses due to ELMs would have on the divertor design of future high power tokamaks such as ITER. Ref.\\cite{1997JNuM..241..182H} reviews what is known about ELMs, with an emphasis on their effect on the scrape-off layer and divertor plasmas. ELM effects have been measured in the ASDEX-U, C-Mod, COMPASS-D, DIlI-D, JET, JFT-2M, JT-60U and TCV tokamaks, and were reported by Ref.\\cite{1997JNuM..241..182H}. ELMs are known to produce super-thermal particle fluxes that are damaging for Tokamak wall materials.  At the beginning of the previous paragraph, interdisciplinary nature of the role of accelerated particles was emphasised. In all of the above mentioned fields of research, the mechanisms that lead to accelerated particles are in fact common, and can be split into three broad classes: (1) Electric field acceleration; (2) Shock acceleration and (3) Stochastic acceleration.  In the electric field acceleration mechanism,  Ref.\\cite{1959PhRv..115..238D} considered dynamics of electrons under the action of two effects: the parallel (to background magnetic field) electric field and friction between electrons and ions. The equation describing electron motion along the magnetic field is $m_e \\dot{v}_d = eE - \\nu_{ei} m_e v_d$, where $v_d$ is the electron drift velocity, $\\nu_{ei}$ is the electron-ion collision frequency and dot denotes time derivative. In the limiting case of $v_d  << v_{thermal}$ and the steady state regime, the equation of motion  allows one to obtain the expression for Spitzer resistivity. In the limit of $v_d  > v_{thermal}$, the steady state solution is not possible. This is because when the right hand side of the equation is positive i.e. when $eE >\\nu_{ei} m_e v_d$, one has electron acceleration. The latter leads to the so-called run-away regime. The acceleration due to the parallel electric field leads to an increase in the electron drift speed $v_d$, in turn, this leads to a decrease in $\\nu_{ei}$ because $\\nu_{ei} \\propto v_d^{-3}$. Therefore, when the electric field exceeds the critical value $E_D=n_e e^3 \\ln \\Lambda/(8 \\pi \\varepsilon_0^2k_BT)$ (so-called Dreicer electric field), the faster electron motion leads to the decrease in electron-ion friction, which in turn results in even faster motion. As a result the run-away regime is reached. For electric fields less than the Dreicer field, the rate of acceleration is less than the rate of collision losses and only a small fraction of the particles can be accelerated into a non-thermal tail of energy $E < LeE_D$ \\cite{2008SSRv..134..207P}, where $L$ is the acceleration lengthscale. Super-Dreicer fields, which seem to be present in many simulations of reconnection \\cite{2006ApJ...644L.145C,2005ApJ...618L.111Z}, accelerate particles at a rate that is faster than the collision or thermalisation time. This can lead to a runaway and an unstable electron distribution which, as shown theoretically, by laboratory experiments and by the simulations \\cite{2006ApJ...644L.145C,2005ApJ...618L.111Z}, will give rise to plasma waves or turbulence \\cite{1970PhRvL..25..706B,1985ApJ...293..584H}.  In shock acceleration mechanism, which is also called first-order Fermi acceleration, the fractional momentum gain, $\\Delta p/p$, per shock crossing is linear in the velocity ratio, $u_{sh}/v$ , where $u_{sh}$ and $v$ are shock and particle speeds respectively. In stochastic acceleration mechanism, which is also called second-order Fermi acceleration, particles of velocity $v$ moving along magnetic field lines with a pitch angle $\\cos\\mu$ undergo random scattering by moving agents with a velocity $u$. Because the head collisions (that are energy gaining) are more probable than trailing collisions (that are energy losing), on average, the particles gain momentum at a rate proportional to $(u / v)^2 D_{\\mu\\mu}$, where $D_{\\mu\\mu}$ is the pitch angle diffusion rate. Note that here dependence on the velocity ratio is quadratic (hence the name second order) \\cite{2008SSRv..134..207P}.  In this work we focus on the electric field acceleration mechanism. In particular when the electric field is provided by the DAWs. In magnetohydrodynamic (MHD) regime Alfven waves (AWs) cannot provide parallel electric fields, therefore cannot accelerate particles via the electric field acceleration mechanism. After it was realised \\cite{1970PhFl...13..440S} that in the kinetic regime (in which plasma is treated as a collection of charged particles, rather than a fluid as in MHD) AWs with parallel electric fields are possible, a significant effort has been undertaken to study these waves. In laboratory plasmas, an externally applied oscillating magnetic field (at a frequency near 1 MHz for typical Tokamak parameters) was found to resonantly mode convert to a kinetic AW, whose perpendicular wavelength is comparable to the ion gyroradius \\cite{1976PhFl...19.1924H}. The kinetic AW, while it propagates into the higher-density side of the plasma after mode conversion, dissipates due to both linear and nonlinear processes and heats the plasma. If a magnetic field of 50 G effective amplitude is applied, approximately 10 MJ m$^{-3}$ of energy can be deposited in 1 sec into the plasma. In space plasmas, surface waves, excited either by a MHD plasma instability or an externally applied impulse, were shown to convert in a resonant mode to a kinetic AW having a wavelength comparable to the ion gyroradius in the direction perpendicular to the magnetic field. The kinetic AW was found to have a component of its electric field in the direction of the ambient magnetic field and can accelerate plasma particles along the field line. A possible relation between this type of acceleration and the formation of auroral arcs was discussed \\cite{1976JGR....81.5083H}. In solar coronal plasmas, it was found that an Alfven surface wave can be excited by 300 s period chromospheric oscillations which then resonantly excite kinetic AWs. Because of their kinetic structure, the latter waves were found to dissipate their energy efficiently \\cite{1978ApJ...226..650I}.  The parallel electric field is also produced when low frequency ($\\omega < \\omega_{ci}$, where $\\omega_{ci}=eB/m_i$ is the ion cyclotron frequency) dispersive AW has a wavelength perpendicular to the background magnetic field comparable to any of the kinetic spatial scales such as: ion gyroradius at electron temperature, $\\rho_s=\\sqrt{k_B T_e/m_i}/ \\omega_{ci}$, ion thermal gyroradius, $\\rho_i=\\sqrt{k_B T_i/m_i}/ \\omega_{ci}$, \\cite{1976JGR....81.5083H} or to electron inertial length $\\lambda_e = c/ \\omega_{pe}$ \\cite{1979JGR....84.7239G}. The central issue in understanding what physical effect(s) can provide the parallel electric field can be understood based on the two-fluid theory of dispersive AWs \\cite{2000SSRv...92..423S}. The electron equation of motion, $\\partial \\vec{v}_e /\\partial t +(\\vec{v}_e \\cdot \\nabla)\\vec{v}_e= -(e/m_e)(\\vec{E}+\\vec{v}_e\\times \\vec{B})-\\nabla \\cdot {P}_e/(m_e n_e)$,  indicates that in the linear regime, $E_\\parallel$ can be supported either by the electron inertia term, $\\partial \\vec{v}_e /\\partial t$, or the electron pressure tensor term, $\\nabla \\cdot {P}_e/(m_e n_e)$. Dispersive Alfven waves are subdivided into Inertial AWs or Kinetic AWs depending on the relation between the plasma $\\beta$ and  electron/ion mass ratio $m_e/m_i$. When $\\beta \\ll m_e/m_i$ (i.e. when Alfven speed is much greater than electron and ion thermal speeds, $V_A \\gg v_{th,i}, v_{th,e}$) dominant mechanism for sustaining $E_\\parallel$ is the parallel electron inertia (that is why such waves are called inertial AWs). When $\\beta \\gg m_e/m_i$, (i.e. when  $V_A \\ll v_{th,i}, v_{th,e}$) then clearly thermal effects become important. Thus, the dominant mechanism for sustaining $E_\\parallel$ is the parallel electron pressure gradient (that is why such waves are called kinetic AWs - kinetic motion of electrons is the source of pressure). In the MHD regime, it has been known \\cite{1983A&A...117..220H} that the phase-mixing mechanism, in which an AW moves along the density inhomogeneity that is transverse to the background magnetic field, can provide a fast dissipation mechanism. Fast in a sense that AW amplitude decrease, as a function of travel distance, $x$, along wave propagation direction, is a much faster function of $x$, $B_\\perp(x)=\\exp[(-\\eta V_A ^{\\prime 2} k^2/6 V_A^3)x^3]$, compared to the usual resistive damping law $B_\\perp(x)= \\exp[(-\\eta k^2/V_A)x]$, where $\\eta$ is the plasma resistivity and $V_A ^{\\prime}$ denotes spatial derivative of the Alfven speed in the transverse to the background magnetic field direction \\cite{1983A&A...117..220H}. The fast dissipation occurs because the phase mixing effect leads to the build-up of sharp gradients in the Alfven wave front, in the transverse to its propagation direction. It has been recently discovered \\cite{2005A&A...435.1105T} in the solar coronal context that when the phase-mixing mechanism is considered in the kinetic regime, electrons can be effectively accelerated by the phase-mixed dispersive AW. It was found \\cite{2005A&A...435.1105T} that when dispersive AW moves in the transversely inhomogeneous plasma, parallel electric field is generated in the regions of plasma inhomogeneity. This electric field can effectively accelerate electrons for a wide range of plasma parameters \\cite{2008PhPl...15k2902T}. It turns out that similar electron acceleration mechanism has been also found in the Earth magnetosphere context \\cite{1999JGR...10422649G,2004AnGeo..22.2081G}. See also Ref.\\cite{2006A&A...449..449M} for the cross-comparison. The difference is that Ref. \\cite{2005A&A...435.1105T,2008PhPl...15k2902T} have studied the case of plasma over-density, transverse to the background magnetic field, whereas Ref.\\cite{1999JGR...10422649G,2004AnGeo..22.2081G,2006A&A...449..449M} considered the case of plasma under-density (a cavity), as dictated by the different applications considered (solar coronal loops and Earth auroral plasma cavities). Ref.\\cite{2008GeoRL..3520101L} studied the kinetic-scale Alfven wave phase mixing at the boundaries of the density cavity, the strongly varying Alfven speed profile above the Earth auroral ionosphere (so-called ionospheric Alfven resonator), which leads to the development of the parallel electric fields that accelerate electrons in the aurora. They performed simulations to study the evolution of these fields including both parallel and perpendicular inhomogeneity. Calculations of Ref.\\cite{2003JGRA..108.8005L}, based on the linear kinetic theory of Alfven waves, also indicate that Landau damping of these waves can efficiently convert the Poynting flux into field-aligned acceleration of electrons. More recently, an analytical theory has been developed \\cite{2011A&A...527A.130B} to explain the numerical simulations presented in Ref.\\cite{2005A&A...435.1105T,2008PhPl...15k2902T}. Also, a possibility of acceleration of electrons by inertial AWs has been explored in a one dimensional, homogeneous plasma case (without phase mixing) \\cite{2009ApJ...693.1494M}. The latter extended and adopted for solar coronal context earlier work of Ref.\\cite{1994JGR....9911095K} who demonstrated the Fermi-like acceleration by inertial Alfven waves of a small proportion of electrons, up to speeds of twice the Alfven speed and argued that this is a plausible scenario for auroral electron acceleration. The effect of the Hall term on the phase mixing of AWs has been studied in Ref.\\cite{2011A&A...525A.155T}.    This work aims to fill the gap in understanding of particle acceleration by dispersive Alfven waves in the transversely inhomogeneous plasma via full kinetic simulation particularly focusing on the effect of polarisation of the waves and different regimes (inertial of kinetic). In particular, we study particle acceleration by the low frequency ($\\omega=0.3\\omega_{ci}$) DAWs, similar to considered in Ref.\\cite{2005A&A...435.1105T,2008PhPl...15k2902T}, but here we focus on the effect of the wave polarisation, left- (L-) and right- (R-) circular and elliptical, in the different regimes inertial ($\\beta < m_e/m_i$) and kinetic ($\\beta > m_e/m_i$). The primary goal is to study the latter effects on the efficiency of particle acceleration and the parallel electric field generation. ",
        "conclusions": "X-ray observations of solar flares indicate that $10^{34}-10^{37}$ per second accelerated electrons are required to produce the observed hard X-ray emission. If the particle acceleration takes place at the loop apex, this implies that for the acceleration volume of $1-10$ Mm$^3$ with the number density of $n=10^{16}$ m$^{-3}$ all 100\\% of electrons need to be accelerated. Given the implausibility of the latter acceleration efficiency, if the flare energy release triggers DAWs, these can rush down the solar coronal loop towards the footpoints and accelerate electrons along the way to chromosphere \\cite{2008ApJ...675.1645F}. In this context and in the light of our initial study Ref.\\cite{2005A&A...435.1105T}, here we study particle acceleration by the low frequency ($\\omega=0.3\\omega_{ci}$) DAWs, which in fact are very similar to usual shear Alfven waves, as far as their dispersion properties are concerned. The DAWs propagate in the transversely inhomogeneous plasma. In which density is increased in the middle, across the magnetic field, by factor of four and temperature is lowered accordingly to keep the structure in pressure balance. Such system mimics a solar coronal loop. DAWs generate parallel electric field only in the density inhomogeneity regions (loop edges). This parallel electric field behaviour coincides with the parallel current  $\\vec J = e (n_i \\vec V_i - n_e \\vec V_e)$. The latter indicates that charge separation (due to different electron and ion inertia) is the cause of the parallel electric field generation. $E_\\parallel$ then accelerates electrons (creates electron beams) along the magnetic field, while $E_\\perp$ mostly heats ions in the transverse direction. In this study we focus on the effect of the wave polarisation, L- and R- circular and elliptical, in the different regimes inertial ($\\beta < m_e/m_i$) and kinetic ($\\beta > m_e/m_i$). Plasma beta is fixed at 0.02. Hence we vary the regimes by setting either $m_e/m_i=1/16=0.0625 > \\beta$ (inertial regime) and $m_e/m_i=1/73.44=0.0136 < \\beta$ (kinetic regime). The mass ratio of $73.44$ is not ideal (it corresponds to $1/25$th or the real mass ratio of $1836$) but it is on the limit of available computational resources as a typical with $m_i/m_e=73.44$ run takes 48 hours on 256 processors (Here we present the results from total of 13 runs). Our main goal was to study the latter effects on the efficiency of particle acceleration and the parallel electric field generation. As a result we have established: (i) In the inertial regime, fraction of accelerated electrons (along the magnetic field), in the density gradient regions is $\\approx 20\\%$ by the time of IC driving of $t_{end}=75\\omega_{ci}^{-1}$ which allows for the $3$ full wavelength to develop and  is increasing to $\\approx 30\\%$ by the time of $t_{end}=300\\omega_{ci}^{-1}$ that allows for the $13$ wavelength to develop. In all considered cases (see Table~\\ref{runs}) ions are heated in the transverse to the magnetic field direction and fraction of heated ions (with velocities above initial average thermal speed) is $\\approx 35 \\%$. (ii) While keeping the power of injected DAWs the same in all considered numerical simulation runs, in the case of right circular, left and right elliptical polarisation DAWs with the electric field in the non-ignorable ($y$) transverse direction exceeding several times that of in the ignorable ($z$) direction produce more pronounced parallel electron beams (with larger maximal electron velocities) and transverse ion beams in the ignorable ($v_z$) direction. In the inertial regime such polarisations yield the fraction of accelerated electrons about  $20\\%$. In the kinetic regime this increases to $35\\%$. (iii) In the case of left circular, left and right elliptical polarisation DAWs with the electric field in the non-ignorable ($y$) transverse direction several times less than that of in the ignorable ($z$) direction produce less energetic parallel electron beams. In the inertial regime such polarisations yield the fraction of accelerated electrons about  $20\\%$. In the kinetic regime this increases to $30\\%$. (iv) The parallel electric field that is generated in the density inhomogeneity regions is independent of the regime (whether kinetic or inertial), i.e. independent of electron-ion mass ratio and stays of the order $0.03\\omega_{pe}c m_e /e$ which for considered solar flaring plasma parameters ($n=10^{16}$ m$^{-3}$ and $T=6\\times10^7$K at the domain edges) exceeds Dreicer electric field $E_D=n_e e^3 \\ln \\Lambda/(8 \\pi \\varepsilon_0^2k_BT)$ (with Coulomb logarithm $\\ln \\Lambda=17.75$) by a factor of $6.4 \\times 10^8$. Although our PIC simulations do not include particle collisions, and thus Dreicer electric field is not directly relevant because this concept is essentially collisional, the fact that the generated parallel electric field exceeds Dreicer value by eight orders of magnitude indicates that the accelerated electrons would be in the run-away regime. The fact that in the present simulations $E_\\parallel$ amplitude is independent of $m_i/m_e$ contradicts with the scaling law ($E_\\parallel \\propto 1/(m_i/m_e)$) found by Ref.\\cite{2007NJPh....9..262T}. The difference can possibly explained by the fact that Ref.\\cite{2007NJPh....9..262T} considered cold plasma approximation i.e. all particles had the same velocity (with the distribution function being a delta-function rather than a Maxwellian). The latter would have prevented the possibility of having wave-particle interactions, which as found in this study are rather important. (v) Electron beam velocity has the phase velocity of the DAW. This can be understood by  Landau damping of DAWs. DAWs give up their energy to electrons via collisionless Landau damping. Because for the DAWs with $\\omega=0.3 \\omega_{ci}$ their phase seed is approximately the Alfven speed, this, combined with the suggestion by Ref.\\cite{2008ApJ...675.1645F} that Alfven speed can be as high as few $\\times 0.1c$, for the realistic mass ratio of $m_i/m_e=1836$, would readily provide electrons with few tens of keV. e.g. for $V_A=0.3c$  the electron energy would be 23 keV. (vi) As we increased the mass ratio from $m_i/m_e=16$ to 73.44 the fraction of accelerated electrons has increased from $20\\%$ to $30-35\\%$ (depending on DAW polarisation). This is because the velocity of the beam has shifted to lower velocity (compare e.g. Fig.3(a) to 9(a)). Since there are {\\it always more electrons with a smaller velocity than higher velocity} in the Maxwellian distribution, for the mass ratio $m_i/m_e=1836$ the fraction of accelerated electrons would be even higher than $35\\%$. (vii) DAWs generate significant density and temperature perturbations that are located in the density gradient regions (solar coronal loop edges). These perturbations of density and temperature should in principle be detectable in  {\\it future} observations because current time cadence of space instruments is well above $\\omega_{ci}^{-1}$.   The overall conclusion is that DAWs, propagating in the transversely inhomogeneous plasma, in the Alfvenic ($\\omega < \\omega_{ci}$) range of the plasma dispersion can effectively accelerate electrons along the magnetic field and heat the ions in the transverse direction. This is bound to have applications in many fields -- Earth magnetosphere and laboratory plasma, in addition to the solar flare context emphasised here."
    },
    "1107/1107.4916_arXiv.txt": {
        "abstract": "In this paper, we compare the observationally derived black hole mass function (BHMF) of luminous ($>10^{45}-10^{46}$ erg/s) broad-line quasars (BLQSOs) at $1<z<4.5$ drawn from the Sloan  Digital Sky Survey (SDSS) presented in Kelly et al. (2010), with models of merger driven BH growth in the context of standard hierarchical structure formation models. In the models, we explore two distinct black hole seeding prescriptions at the highest redshifts: \"light seeds\" - remnants of Population III stars and \"massive seeds\" that form from the direct collapse of pre-galactic disks. The subsequent merger triggered mass build-up of the black hole population is tracked over cosmic time under the assumption of a fixed accretion rate as well as rates drawn from the distribution derived by Merloni \\& Heinz (2008). Four model snapshots at $z=1.25$, $z=2$, $z=3.25$, $z=4.25$ are compared to the SDSS derived BHMFs of BLQSOs. We find that the light seed models fall short of reproducing the observationally derived mass function of BLQSOs  at $M_{\\rm BH}>10^9 \\msun$ throughout the redshift range; the massive seed models with a fixed accretion rate of $0.3\\,Edd$, or with accretion rates drawn from the Merloni \\&  Heinz distribution provide the best fit to the current observational data at $z>2$, although they overestimate the high-mass end of the mass function at lower redshifts. At low redshifts, a drastic drop in the accretion rate is observed and this is explained as arising due to the diminished gas supply available due to consumption by star formation or changes in the geometry of the inner feeding regions. Therefore, the over-estimate at the high mass end of the black hole mass function for the massive seed models be easily be modified, as the accretion rate is likely significantly lower at these epochs than what we assume. For the Merloni \\& Heinz (2008) model, examining the Eddington ratio distributions $f_{Edd}$, we find that they are almost uniformly sampled from $f_{Edd}=10^{-2}-1$ at $z\\simeq 1$, while at high redshift, current observations suggest accretion rates close to Eddington, if not mildly super-Eddington, at least for these extremely luminous quasars. Our key findings are that the duty cycle of SMBHs powering BLQSOs increases with increasing redshift for all models and models with Pop III remnants as black hole seeds are unable to fit the observationally derived BHMFs for BLQSOs, lending strong support for the massive seeding model. ",
        "introduction": "Demography of local galaxies suggests the that most galaxies harbour a quiescent super-massive black hole (SMBH) in their nucleus at the present time and the mass of the hosted SMBH is correlated with properties of the host bulge. Observational evidence supports the existence of a strong connection between the growth of the central SMBH and stellar assembly of the host spheroid (Tremaine et al. 2002; Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Marconi \\& Hunt 2003; H\\'aring \\& Rix 2004; G\\\"ultekin et al. 2009) and possibly the host halo in nearby galaxies (Ferrarese 2002). This strong correlation is suggestive of co-eval growth of the BH and the stellar component likely via regulation of the gas supply in galactic nuclei from the earliest times (Silk \\& Rees 1998; Kauffmann \\& Haehnelt 2000; Fabian 2002; King 2003; Thompson, Quataert \\& Murray 2005; Natarajan \\& Treister 2009; Treister et al. 2011). Meanwhile, our observational understanding of accreting black holes and their properties at high redshifts $z>6$ is growing (see Fan et al. 2010 for an overview). Black holes appear to be copiously accreting and in place before the universe was a billion years old as evidenced by the recent detection of the quasar at $z=7.085$ in the UKIDSS survey reported by Mortlock et al. (2011). Our access to the earliest redshifts is of course restricted to the brightest candidates, optical, and deep X-ray surveys are needed to uncover the potentially highly obscured population at these epochs (Treister et al. 2011).  Simultaneously our theoretical understanding of the assembly history of black holes over cosmic time has been rapidly improving (see recent reviews by Volonteri 2010 and Natarajan 2011). With the increasing availability of observations from earlier and earlier epochs, models can be effectively tested at extremely high redshifts. Recently, a more complete census of the accreting black hole population at $1 < z < 4.5$ has become available enabling the estimate of the mass function (MF) of these black holes (Kelly et al. 2010, K10 hereafter). In this paper, we explore a range of theoretically viable models to fit this additional and new data-set.  Observations currently provide strong constraints locally, and hints at high redshifts (Mortlock et al. 2011; Treister et al. 2011). With this additional anchor during inter-mediate epochs ($z\\,\\sim\\,1-4$), discrimination between models finally starts becoming possible as we demonstrate here. In particular, we are able to distinguish between seed black hole formation scenarios. This paper is the third in a series (Volonteri, Lodato \\&Natarajan 2008, and Volonteri \\& Natarajan 2009, VLN08 and VN09 respectively hereafter) where we compare the late time evolution of SMBH seeds, focusing on the difference between a population of more numerous, but smaller mass BHs as seeds, versus a population of rarer, but more massive initial seeds. The outline of this paper is as follows: in Section 1, we briefly review the new observations and BHMFs derived therefrom; Section 2 provides the context of the theoretical black hole growth models followed by a description of how BHMFs are derived in Section 3; and finally we present the results of comparing observations to the models and the derived conclusions. ",
        "conclusions": "In what follows, we present a detailed comparison of the data with our three models, namely the {\\bf PopIII-Edd}: initial seeds from Pop III remnants with accretion assumed at all times at the Eddington rate; {\\bf Massive-MH}: wherein the initial massive seeds have accretion rates drawn from the distribution determined by MH08; {\\bf Massive-subEdd}: initial massive seeds with accretion assumed at all times to be at $0.3 \\times$ the Eddington rate. For all three models, we compare our derived MF of BLQSOs with that estimated by K10 and the MF of all SMBHs (the entire population that includes actively accreting + inactive black holes) with the MF predicted by MH08. Although the latter is also a theoretical MF, it is derived under completely different assumptions, as the MH08 model starts with SMBHs today and evolves them back in time via a continuity equation. We therefore consider this an interesting case of model confronting model (see also Volonteri \\& Begelman 2010).  We compare our derived MFs for BLQSOs to K10 and our MFs for the entire BH population (active + inactive) to MH08 in Figures~1-4. We find the following: \\begin{itemize}  \\item In the higher redshift slices at $z=4.25$ and $z=3.25$, independent of our models the MF derived in MH08 for all SMBHs is incompatible with that for the subset of BLQSOs from K10. In these instances the abundance of BLQSOs from K10 is in excess of that of all SMBHs from MH08 models. It is only in the lowest redshift bin at $z=1.25$ that the MFs of BLQSOs from K10 is lower than that of all SMBHs derived from the MH08 models. It is therefore clear that the Eddington ratio distributions derived from K10 and from  MH08 are most likely discrepant.    \\item At the highest redshift slice $z=4.25$, the light seed model {\\bf PopIII-Edd}) shown in the top panel fails to grow SMBHs massive enough (more massive than $\\sim 10^{8.5} M_{\\odot}$) to populate the high-mass end of the MF of BLQSOs of K10. The discrepancy is roughly two orders of magnitude as clearly seen in the upper panel of Figure~1. Since the Eddington ratio was assumed to be unity to evolve this model, a deficit at the high mass end cannot be compensated for by altering the accretion physics any further.  Alternatively, if SMBHs can become ``over massive\" at higher z (i.e., if they climb well above the standard $M_{\\rm BH}-\\sigma$ relation) and have phases of super-Eddington accretion (Volonteri \\& Rees 2005), then enough SMBHs at the high mass end might be produced. We note that assuming the MH08 accretion model in conjunction with Pop III remnants (this model is not plotted in our figures) makes the discrepancy at the high mass end much worse, as the accretion rate for low mass black holes ($<1000\\, \\msun$) in the MH08 model peaks at sub-Eddington rates at all redshifts.  \\item Both the massive seed models ({\\bf Massive-subEdd} and {\\bf Massive-MH}) also under-produce the K10 MFs at the highest redshift slice as seen clearly by looking at the middle and bottom panels in $z = 4.25$;  \\item The MF of all existing SMBHs (active + inactive) from the {\\bf PopIII-Edd}, {\\bf Massive-subEdd} and {\\bf Massive-MH} models overestimates MH08's MF for all SMBHs at all masses $>3\\times 10^7\\, \\msun$ at $z=3.25$, while there is reasonably good agreement with K10 for the BLQSO MF for the {\\bf Massive-MH} model.  \\item At lower redshifts ($z=1.25$ and $z=2$) both massive seed models start to overestimate the MF of BLQSOs  by K10 at $M_{\\rm BH}>10^9\\,\\msun$.  This is not unexpected, as we have not implemented a cut-off to mimic the depletion of available gas in massive galaxies.  \\item All massive seeds models ({\\bf Massive-subEdd} and {\\bf Massive-MH}) overestimate the MF of all SMBHs (active + inactive) from MH08 at $M_{\\rm BH}>10^{8.5}\\, \\msun$. At $z<2$ we  cannot clearly assess whether MH08 might underestimate the total MF. We are, however, inclined to attribute the discrepancy to our models,  for the same reason that we likely overestimate the MF of quasars as we do not take into account changes in the gas inventory at late times in the universe. The {\\bf PopIII-Edd} model gets instead into better alignment at lower and lower redshifts. At $z = 2$ these MFs begin to agree up to $10^8\\,M_{\\odot}$ and at $z = 1.25$ they fall into excellent agreement up to $10^9\\,M_{\\odot}$.   \\item For model {\\bf Massive-MH} we can estimate the typical Eddington ratio of SMBHs powering BLQSOs. We find that most luminous quasars at the highest redshift have accretion rates close to or even slightly larger than the Eddington rate (the fraction of super-Eddington accreting SMBHs is $\\simeq$ 60\\% at $z=3$; $\\simeq$ 40\\% at $z=2$; $\\simeq$ 20\\% at $z=1$). At lower redshift, the flux limit of the SDSS survey corresponds to lower luminosities, and less powerful accretors are selected. In general, we find that the accretion rates in the models overestimate the accretion rates derived by K10.  \\item We estimate the duty cycle, defined as the fraction of SMBHs that are active and detectable as BLQSOs, and show the results for the three models ({\\bf PopIII-Edd}, {\\bf Massive-subEdd}, {\\bf Massive-MH}) in Figure~\\ref{dc} at $z=1$, $z=2$, $z=3$, $z=4$ (dark to light shade of colours). Comparing the theoretical duty cycle with the duty cycle estimated by K10 at $z=1$, we find that model {\\bf PopIII-Edd} is consistent at all masses; while models {\\bf Massive-MH} and {\\bf Massive-Edd} overestimate the duty cycle at $M_{\\rm BH}>10^9\\, \\msun$ (as expected from the MF shown in Figures 3 and 4.).  \\item We find that the duty cycle increases with increasing redshift for all our models, although more strongly for the {\\bf Massive-subEdd} model. This is due to the fact that this model is the one with the lowest average accretion rate ($f_{\\rm Edd}=0.3$, versus  $f_{\\rm Edd}=1$ for the {\\bf PopIII-Edd} and  $f_{\\rm Edd}\\gsim 1$ in {\\bf Massive-MH} at $z>3$, therefore, from Equation~\\ref{tagn} the BHs in the {\\bf Massive-subEdd} model need to accrete for a longer time in order to reach the final mass set by the $M_{\\rm BH}-\\sigma$ relation.  \\end{itemize}  In conclusion, the implications for seeding models are that light seeds always under-produce the bright end of the BLQSO LF for high black hole masses. This is a serious issue as they are assumed to be already accreting at the optimal Eddington rate.  Moreover, according to recent work (Turk et al. 2009; Greif et al. 2011; Davis et al. 2011) it appears now that (i) the IMF of the first stars may not be biased high as fragmentation might occur during the formation process and (ii) due to turbulence the formation of the first stars might get delayed. So the {\\bf PopIII-Edd} model aside from being unable to match the black hole MF derived from BLQSOs from the SDSS as shown here, might in fact not be the most efficient channel to produce seeds. It is clear that a channel to form more massive black hole seeds at high redshifts is required to start matching observations of the bright end (massive end of the black hole MF) from the earliest to intermediate redshifts ($z \\sim 1$), and not only at the highest redshifts as previously noted (e.g., Haiman 2004; Yoo \\& Miralda-Escud{\\'e} 2004; Shapiro 2005;  Volonteri \\& Rees 2005, 2006). {\\bf Therefore, low-mass SMBH seeds such as Pop III star remnants require exceptional growth conditions not only to explain the existence of rare objects such as $z>6$ quasars, but also to explain the unexceptional population of luminous quasars at $z\\simeq 2-3$.}  Over-predicting the BHMF at the high mass end is typically not a serious problem as lowering the Eddington rate would fix that.  For instance, Natarajan \\& Treister (2009) propose the existence of an upper mass limit to BHs at all epochs, arising from the shutting off of accretion due to self-regulation of the gas supply in the galactic nuclei. According to their estimate this upper limit is given by: \\begin{eqnarray} M_{\\rm BH}  \\sim 5 \\times 10^9\\, ({\\frac{\\sigma}{350\\,{\\rm km\\,s}^{-1}}})^5\\,\\msun. \\end{eqnarray} K10 estimate an upper mass limit at $z \\sim 1$ for a black hole hosting a BLQSO to be $M_{\\rm BH} \\sim 3 \\times 10^{10}\\,\\msun$. Utilizing the limit proposed by Natarajan \\& Treister (2009), if a BH shuts off at the end of this BLQSO episode at $z \\sim 1$, it is likely hosted in a galaxy with a velocity dispersion of $\\sigma \\sim$ 440 km/s, which corresponds to the typical velocity dispersion of the brightest central galaxy in a cluster. This fits in nicely with the view of down-sizing wherein accretion is noted to shut off at $z \\sim 1$ in these BCGs in clusters.  \\begin{figure} \\includegraphics[width= \\columnwidth]{fig4.eps} \\caption{The duty cycle of BLQSOs as a function of SMBH mass. The models shown here are again in the same order as before: ${\\bf PopIII-Edd}$ (uppermost panel), ${\\bf Massive-MH}$ (middle panel), ${\\bf Massive-subEdd}$ (bottom panel). In each panel, we show the duty cycle at $z=4$, $z=3$, $z=2$, and $z=1$ (from lightest to darkest shade respectively). The square symbols show the lower bound on the duty cycle at $z=1$ from K10.} \\label{dc} \\end{figure}"
    },
    "1107/1107.5910_arXiv.txt": {
        "abstract": "{ The standard Friedmann model of cosmology is based on the Copernican Principle, i.e. the assumption of a homogeneous background on which structure forms via perturbations. Homogeneity underpins both general relativistic and modified gravity models and is central to the way in which we interpret observations of the CMB and the galaxy distribution. It is therefore important to probe homogeneity via observations. We describe a test based on the fossil record of distant galaxies: if we can reconstruct key intrinsic properties of galaxies as functions of proper time along their worldlines, we can compare such properties at the same proper time for our galaxy and others.  We achieve this by computing the lookback time using radial Baryon Acoustic Oscillations,  and the time along galaxy world line using stellar physics, allowing us to probe homogeneity, in principle anywhere inside the past light cone.  Agreement in the results would be an important consistency test -- although it would not in itself prove homogeneity. Any significant deviation in the results however would signal a breakdown of homogeneity. } ",
        "introduction": "The standard model of the Universe is homogeneous, with structure formation described via perturbations on a fixed background. Within general relativity (GR), it then follows that the acceleration of the Universe is driven by dark energy. The homogeneous standard model is a simple, predictive model that successfully accommodates all observations up to now. However, we should probe the foundations of this model as far as possible. This is further motivated by the fact that there is still no satisfactory description of the dark energy that is central to the model. We can probe the assumption that GR holds on cosmological scales, by investigating modified gravity theories and by devising consistency tests of GR. This probe is only effective if we assume homogeneity. Alternatively, we can probe the assumption of homogeneity itself \\cite{Maa11,Ellis:2011hk,ClaMaa10}.  Homogeneity is {\\em not} established by observations of the CMB and the galaxy distribution -- {\\em we cannot directly observe homogeneity,} since we observe down the past lightcone, recording properties on 2-spheres of constant redshift and not on spatial surfaces that intersect that lightcone. What these observations can directly probe is isotropy. In order to link isotropy to homogeneity, we have to assume the Copernican Principle, i.e. that we are not at a special position in the Universe. The Copernican Principle is not observationally based; it is an expression of the intrinsic limitation of observations from one spacetime location.  We cannot observationally prove homogeneity, but we can develop consistency tests that would uncover deviations from homogeneity if these deviations existed. In other words, we can test for consistency of the Copernican Principle and for violations of the Copernican Principle.  If we find no violation, then the indirect evidence for homogeneity is strengthened. However,  any single significant violation of the Copernican Principle would disprove homogeneity.   Here we propose a new test of homogeneity, based on intrinsic properties of galaxy evolution. Homogeneity means that fundamental properties, such as star formation histories, should be uniform -- this is the `fossil record' implicitly carried by galaxies via their evolution. By observing distant galaxies, and then tracing our own galactic history back to the proper time of emission, we can check for consistency between properties of our galaxy and distant galaxies.  This is closely related to the `cosmic chronometer' project \\cite{Stern:2009ep,Crawford:2010rg,Moresco:2010wh}, which uses galaxy spectral properties to find $H(z)$ in a Friedmann model (see also \\cite{Carson:2010sq}). Here we do not assume Friedmann geometry but rather test for it. It is also related to earlier work which attempted to prove homogeneity from uniformity of the thermal history of galaxies \\cite{bonell86}. However it turns out that there are inhomogeneous models with uniform thermal histories \\cite{bonell86}, and the proposal fails as a proof of homogeneity. It does survive however as a probe of homogeneity -- i.e., violations of uniformity would indicate a breakdown of the Copernican Principle and thus of homogeneity. The question is how to apply the idea in practice.  Before describing our proposal, we briefly review other tests of the Copernican Principle and homogeneity.    \\begin{itemize} \\item  The geometric null test of \\cite{CBL} is based on the tight relation between curvature and expansion history in a Friedmann spacetime. This relation is independent of dark energy and other contributions to the energy-momentum tensor (and also independent of the field equations). The quantity \\begin{eqnarray} {\\cal K}(z) &:=& H_0^2(1+z)^4 + H^2(z)\\left[(1+z)^2\\left\\{D_L(z)D_L''(z)- D_L^{\\prime 2}(z)\\right\\}+D_L^2(z) \\right] \\nonumber\\\\ &&~ + (1+z)H(z) H'(z)D_L(z)\\left[(1+z) D_L'(z)-D_L(z) \\right]  \\label{cz test} \\end{eqnarray} is identically zero at all redshifts in a Friedmann spacetime with arbitrary content and curvature. It follows that \\be {\\cal K}(z) ~~\\mbox{significantly different from 0} ~~\\Rightarrow % ~~ \\mbox{violation of homogeneity.} \\ee (How to implement the test is discussed in \\cite{Shafieloo:2009hi}.)  \\item The baryon acoustic oscillation (BAO) feature in the galaxy distribution provides a further geometric test of homogeneity % \\cite{CBL,Clarkson:2010ej,Maa11}. Future large-volume surveys will allow the detection of the BAO scale in both radial and transverse directions. These scales are equal in a Friedmann background, so that \\be {\\mbox{radial BAO} \\over \\mbox{transverse BAO}}-1  ~~\\mbox{significantly different from 0} ~~\\Rightarrow ~~ \\mbox{violation of homogeneity.} \\ee  \\item  Galaxy clusters with their hot ionized intra-cluster gas, act via scattering of CMB photons like giant mirrors that carry information about the last scattering surface seen by the cluster -- and therefore serve as indirect probes inside our past lightcone. The Sunyaev-Zeldovich (SZ) effect on the CMB temperature anisotropies probes the monopole (thermal SZ) and dipole (kinetic SZ) seen by the cluster in a perturbed Friedmann model. Thus \\cite{goodman,Maa11} \\be \\mbox{Non-perturbative thermal or kinetic SZ temperature effect}  ~~\\Rightarrow ~~ \\mbox{violation of homogeneity.} \\ee (In related work, the SZ effect has been used to place constraints on inhomogeneous void models % \\cite{CS,Zhang:2010fa,Moss:2011ze}.)  The SZ effect on CMB {\\em polarization} provides a new and different probe of the monopole and dipole, and in addition probes the quadrupole and octupole seen by the cluster, in a perturbed Friedmann model. Thus \\cite{Maa11}: \\be \\mbox{Non-perturbative SZ polarization effects}  ~~\\Rightarrow ~~ \\mbox{violation of homogeneity.} \\ee  \\end{itemize} Other tests that may become possible with future observations include: a geometric null test based on the time drift of the cosmological redshift \\cite{UCE}; massive neutrinos propagate inside our past lightcone and provide in principle a probe of deviations from homogeneity \\cite{Jia:2008ti}, if the cosmic neutrino background can be observed.   We now turn to our proposed probe of the Copernican Principle, based on the fossil record carried by galaxies, which is also an indirect probe of conditions inside our past lightcone. First we must deal with the problem of lookback time. ",
        "conclusions": "\\label{discon}    Although we cannot directly observe homogeneity, we can test the Copernican Principle at the foundation of homogeneity, using observations that carry information from inside our past lightcone. Previously this has been discussed via consistency relations between distances and expansion rates, using supernova and BAO data, and via Sunyaev-Zeldovich effects from galaxy clusters on the CMB temperature and polarization. Here we have proposed a new test of homogeneity based on the fossil record of distant galaxies. We make observations at constant time slices within our past light cone by computing the lookback time to a distant galaxy and adding it to the proper time along the galaxy past worldline.  The  lookback time is computed without assuming Friedmann geometry, using standard rulers (the radial BAO); the proper time on the galaxy worldline is effectively computed using the fossil record, dependent on the physics of stellar evolution but not on a Friedmann geometry.  As remarked, this test has the advantage that it allows one to look {\\em inside} our past light cone, in principle everywhere inside.  Naturally, we cannot make any global statements about homogeneity, as we are inevitably restricted to the region on or within our past light cone.  The method is interesting in the sense that once the lookback time has been computed, then the fossil record comparisons should be limited only by observational error and sample variance, and if fossil record observations are not consistent in different locations then homogeneity does not apply.  The converse is that any failure to find inhomogeneity strengthens the case for the Copernican Principle.    \\[ \\]{\\bf Acknowledgments:} We thank Chris Clarkson,  Catherine Cress, George Ellis, Licia Verde and Christian Wagner for discussions. RM is supported by a South African SKA Research Chair, by the UK Science \\& Technology Facilities Council (grant no. ST/H002774/1) and by a NRF (South Africa)/ Royal Society (UK) exchange grant."
    },
    "1107/1107.0704_arXiv.txt": {
        "abstract": "{ We report on very-long-baseline-interferometry (VLBI) monitoring observations of the low-luminosity active galactic nucleus (LLAGN) in the galaxy M\\,81 at the frequencies of 1.7, 2.3, 5.0, and 8.4\\,GHz. The observations reported here are phase-referenced to the supernova SN\\,1993J (located in the same galaxy) and cover from late 1993 to late 2005. The large amount of available observations allows us to study the stability of the AGN position in the frame of its host galaxy at different frequencies and chromatic effects in the jet morphology, together with their time evolution. The source consists at all frequencies of a slightly resolved core and a small jet extension towards the northeast direction (position angle of $\\sim$65 degrees) in agreement with previous publications. We find that the position of the intensity peak in the images at 8.4\\,GHz is very stable in the galactic frame of M\\,81 (proper motion upper limit about 10$\\mu$as per year). We confirm previous reports that the peaks at all frequencies are systematically shifted among them, possibly due to opacity effects in the jet as predicted by the standard relativistic jet model. We use this model, under plausible assumptions, to estimate the magnetic field in the jet close to the jet base and the mass of the central black hole. We obtain a black-hole mass of $\\sim$$2\\times10^7$\\,M$_{\\odot}$, comparable to estimates previously reported using different approaches, but the magnetic fields obtained are $10^3-10^4$ times lower than previous estimates. We find that the positions of the cores at 1.7, 2.3, and 5.0\\,GHz are less stable than that at 8.4\\,GHz and evolve systematically, shifting southward at a rate of several tens of $\\mu$as per year. The evolution in the jet orientation seems to be related to changes in the inclination of the cores at all frequencies. These results can be interpreted as due to a precessing jet. The evolving jet orientation also seems to be related to a flare in the peak flux densities at 5.0 and 8.4\\,GHz, which lasts $\\sim$4 years (from mid 1997 to mid 2001). An increase in the accretion rate of the black hole, and its correlation with the jet luminosity via the disk-jet connection model, seems insufficient to explain this long flare and the simultaneous evolution in the jet orientation. A continued monitoring of the flux density and the jet structure evolution in this LLAGN will be necessary to further confirm our jet precession model.} ",
        "introduction": "\\label{I}  The galaxy M\\,81 (NGC\\,3031), located at a distance of $3.63\\pm0.34$\\,Mpc (Freedman et al. \\cite{Freedman1994}), hosts the next closest active galactic nucleus (AGN) after Centaurus A. Its high declination (J2000.0 coordinates $\\alpha = 09^\\mathrm{h}\\,55^\\mathrm{m}\\,33.173^\\mathrm{s}$ and $\\delta = 69^{\\circ}\\,03'\\,55.062''$) also makes M\\,81 a perfect target for global VLBI observations using the most sensitive radio telescopes in the Northern Hemisphere.  The flux density of the AGN in M\\,81 (hereafter, M\\,81*), has relatively large variability, from X-rays (Ishisaki et al. \\cite{Ishisaki1996}; Page et al. \\cite{Page2004}) to radio wavelengths (Ho et al. \\cite{Ho1999} and Bietenholz, Bartel \\& Rupen \\cite{Bietenholz2000}, at centimeter wavelengths; Sch\\\"odel et al. \\cite{Schodel2008}, at millimeter wavelengths). The variability time scale runs from years to days, indicating a contribution from a very compact source to the emission at all frequencies. Indeed, even from VLBI observations at 43\\,GHz (Ros \\& P\\'erez-Torres \\cite{Ros2008}) and space-VLBI observations at 5\\,GHz (Bartel \\& Bietenholz \\cite{Bartel2000}) the emitting structure could not be fully resolved. The broadband luminosity of M\\,81* is relatively weak (see, e.g., Ho et al. \\cite{Ho1996}) from radio (around 100\\,mJy, i.e., on the order of $10^{37}$\\,erg\\,s$^{-1}$) to the optical and X-rays ($\\sim10^{40}$\\,erg\\,s$^{-1}$). Therefore, it is classified as a low-luminosity AGN (LLAGN).  The accretion mechanisms in LLAGNs are not clear. Since relativistic X-ray lines are not clearly detected in these objects (e.g., Reynolds et al. \\cite{Reynolds2009}), radial flows or advection mechanisms might dominate the accretion, in contrast to the more luminuos AGN. In any case, if the X-ray emission traces the innermost inflow from the disk to the central engine and the radio emission traces the jet power, there should be a relationship between luminosities in both X-rays and radio for all AGNs (i.e., including LLAGNs), which should also be related to the accretion rate. This is the basis of the {\\em fundamental-plane} model of black-hole accretion (e.g., Merloni, Heinz, \\& di Matteo \\cite{Merloni2003}). However, the parametric space of very-low accretion flows in AGNs (e.g., LLAGN) is not fully characterized observationally. Although Markoff et al. (\\cite{Markoff2008}) and Miller et al. (\\cite{Miller2010}) were able to fit average broadband spectra of M\\,81* in the frame of the fundamental-plane model of disk-jet connection, no clear correlation between the variabilities in the X-rays and radio was observed unlike in microquasars (Mirabel et al. \\cite{Mirabel1998}) and AGNs (Marscher et al. \\cite{Marscher2002}). In the present paper, we report on a very long flare of M\\,81* at radio frequencies, beginning around mid 1997, which lasted nearly four years. As shown in this paper, an intriguing result related to this flare is that it seems to be related to the evolution in the jet orientation, thus making it difficult explain in the frame of the disk-jet connection between accretion flow and jet power.  At radio frequencies, M\\,81* has a slightly inverted spectrum with an spectral index $\\alpha = +0.3$ ($S\\propto\\nu^{+\\alpha}$), up to a turnover frequency of $\\sim$200\\,GHz (Reuter \\& Lesch \\cite{Reuter1996}). If due to synchrotron self-absorption in the jet, such a high turnover frequency, suggests a very dense plasma and a high magnetic field, and/or a rather flat energy distribution of the relativistic synchrotron-emitting particles (see, e.g., Sect. 4.1 of Reuter \\& Lesch \\cite{Reuter1996}). {In addition, it is worth noticing that} the spectral behavior of M\\,81* is similar to that of the central engine in our Galaxy, Sgr\\,A* (Duschl \\& Lesch \\cite{Duschl1994}), thus suggesting that similar physical mechanisms should be present in the acceleration of the relativistic particles and their emission at radio frequencies. In addition, the mass of its associated black hole ($7^{+2}_{-1}\\times10^{7}$\\,M$_{\\odot}$, Devereux et al. \\cite{Devereux2003}; $5.5^{+3.6}_{2.0}\\times10^{7}$\\,M$_{\\odot}$, Schorr M\\\"uller et al. \\cite{Schorr2011}) makes M\\,81* an interesting intermediate object between Sgr A* and the more powerful AGN.   The explosion of the radio-luminous supernova SN\\,1993J in M\\,81, which happend around 28 March 1993 (e.g., Weiler et al. \\cite{Weiler2007}), triggered an intense campaign of VLBI observations in which M\\,81* was selected as the phase calibrator of the SN\\,1993J visibilities. These observations allow us to study the evolution in the structure and position of M\\,81*, in the frame of its host galaxy and at different frequencies and times. Results for some of the M\\,81* observations at 8.4\\,GHz between years 1993 and 1997 are reported in Bietenholz, Bartel \\& Rupen (\\cite{Bietenholz2000}), and results for the astrometry analysis of a subset of observations at all frequencies between years 1993 and 2002 were reported in Bietenholz et al. (\\cite{Bietenholz2004}). These authors conclude that the proper motion of M\\,81*, relative to the shell center of SN\\,1993J, was consistent with zero up to the precision level of the observations ($11.4 \\pm 9.3$\\,$\\mu$as\\,yr$^{-1}$). The contributions of galactic rotation and/or any peculiar motion of the SN\\,1993J progenitor were also discarded to contribute to a detectable proper motion of M\\,81* relative to the supernova. These authors also report a frequency-dependent shift of the peak of brightness of M\\,81*, which they interpret as opacity effects in the jet of the AGN, in agreement with the standard relativistic jet model (Blandford \\& K\\\"onigl \\cite{Blandford}).  In this paper, we report on the analysis of all the VLBI observations of M\\,81* taken in the phase-referencing observing campaign of SN\\,1993J, from late 1993 to late 2005. In Sect. \\ref{ObsSecVLBI} we describe our observations. In Sects. \\ref{StrucSec} and \\ref{ResultsSecII} we report the main results obtained (source morphology and multi-frequency astrometry). In Sect. \\ref{DiscussionSec}, we analyze the astrometry results in terms of the standard jet interaction model (Blandford \\& K\\\"onigl \\cite{Blandford}), and in the last section we discuss the detection of a changing jet orientation, which we interpret as due to jet precession. ",
        "conclusions": "\\label{ConclusionsSec}  We have reported on a set of global VLBI observations of the AGN in M81 (M\\,81*), performed between years 1993 and 2005 at the frequencies of 1.7, 2.3, 5.0, and 8.4\\,GHz. We studied the morphological evolution of the source and its differential astrometry with respect to SN\\,1993J (a stable position reference in the galactic frame of M\\,81).  The source consists, at all frequencies, of a slightly resolved core component plus a weak jet extension in the northeast direction (position angle of about 65 degrees). The core component has an increasing size with decreasing frequency (from $\\sim0.5$\\,mas at 8.4\\,GHz to $\\sim2.2$\\,mas at 1.7\\,GHz). At 1.7\\,GHz we also find a more extended emission located at$\\sim15$\\,mas from the core and at a position angle of $\\sim65$ degrees. This extended emission hints at a rather straight jet in M\\,81* at least up to 0.26\\,pc (1.1\\,pc) of projected (unprojected) distance to the core. The core components at all frequencies are elongated roughly in the direction of the jet extension and show evidence of an increasing position angle.  The evolution of the peak flux density of the core shows a clear long-term trend at 5.0 and 8.4\\,GHz, with an increase starting around June 1997, a maximum reached around July 1998, and a decrease ending around June 2001. The flux-density variability during this long-term flare is about 100\\% with respect to that in the quiescent stage.  For the astrometric study, we find a shift in the brightness peak of M\\,81* towards the southwest direction with increasing frequency, as previously reported in Bietenholz, Bartel, \\& Rupen (\\cite{Bietenholz2004}). We compared these shifts with the predictions of the relativistic jet model, assuming a power-law behavior for the decay in the magnetic field and the particle density distribution along the jet. We used the Lobanov (\\cite{Lobanov98}) model to estimate the position of the central engine of M\\,81* (i.e., the origin of the jet), relative to the average position of the brightness peaks of our VLBI images. The estimated position of the jet origin is in excellent agreement with the one reported in previous publications using different approaches (Bietenholz, Bartel, \\& Rupen \\cite{Bietenholz2000}, \\cite{Bietenholz2004}). We also used the model of Lobanov (\\cite{Lobanov98}) to estimate the magnetic field in the jet and the mass of the central black hole. We obtained magnetic fields of a few milli-Gauss at distances of a fraction of a parsec from the central engine of the AGN. These values are much lower than those needed to fit the (self-absorption dominated) inverted spectrum of M\\,81*. However, we obtained a black-hole mass of $\\sim$$2\\times10^7$\\,$M_{\\odot}$, comparable to previously reported estimates using different approaches.  From the spectral-index images of M\\,81*, we find a region with inverted spectrum (typical values of $\\alpha$ between 0.5 and 1.5) around the intensity peak at each frequency. The spectrum becomes steeper along the jet. This behavior, together with the multiwavelength astrometric results, agrees with the relativistic jet model.  The position of the peak intensity at 8.4\\,GHz is very stable in the frame of the host galaxy, with a proper motion consistent with zero to better than 10\\,$\\mu$as\\,year$^{-1}$. The positions of the peaks at the other frequencies are less stable, and they systematically evolve towards the south (i.e., larger position angles) at a rate of 10, 76, and 34\\,$\\mu$as\\,year$^{-1}$ at 1.7, 2.3, and 5.0\\,GHz, respectively. We interpret these results as strong evidence of jet precession towards the south.  There is also a relation between the evolution in the relative astrometry among the peaks at different frequencies and the evolving inclination of the cores, as fitted with elliptical Gaussians. This relation lends further support to an evolving orientation of the overall jet of M\\,81*. The long-term variability found in the peak flux densities at 5.0 and 8.4\\,GHz also seems related to the evolving jet orientation. In such a case, if the changing orientation of the jet is periodic (i.e., the jet is precessing, so that its position angle will eventually turn northward after reaching its maximum value, around 70 degrees, and then turn back eastward), we predict that the flux density of M\\,81* at the higher frequencies will eventually increase and follow a light curve similar to our observations by the next 10 years. Thus, a continued monitoring of the flux density and jet structure evolution of this LLAGN will be necessary to further confirm (and, eventually, better constrain) our precession model.  An increase in the accretion rate of the central engine and its later correlation with the jet particle density seem unlikely to explain the light curve of Fig. \\ref{CoreMod1} (left), since it would not explain the simultaneous change in the jet orientation. Therefore, one would be tempted to think of a correlation between the flare and the evolving jet orientation, which would, in principle, imply the periodic generation of similar long-term flares in the future, as in the case of B0605-085 (Kudryavtseva et al. \\cite{Kudriat2011})."
    },
    "1107/1107.5317_arXiv.txt": {
        "abstract": "The temporal heating and subsequent cooling of the crusts of transiently accreting neutron stars carries unique information about their structure and a variety of nuclear reaction processes. We report on a new \\chan\\ Director's Discretionary Time observation of the globular cluster Terzan 5, aimed to monitor the transiently accreting 11-Hz X-ray pulsar \\igr\\ after the cessation of its recent 10-week long accretion outburst. During the observation, which was performed $\\simeq125$ days into quiescence, the source displays a thermal spectrum that fits to a neutron star atmosphere model with a temperature for an observer at infinity of $kT^{\\infty} \\simeq 92$~eV. This is $\\simeq10\\%$ lower than found $\\simeq75$ days earlier, yet $\\simeq 20 \\%$ higher than the quiescent base level measured prior to the recent outburst. This can be interpreted as cooling of the accretion-heated neutron star crust, and implies that crust cooling is observable after short accretion episodes. Comparison with neutron star thermal evolution simulations indicates that substantial heat must be released at shallow depth inside the neutron star, which is not accounted for in current nuclear heating models. ",
        "introduction": "\\label{sec:intro} Neutron stars are the densest directly observable stellar objects in our Universe and constitute ideal astrophysical laboratories to study matter under extreme physical conditions. The outer layer of a neutron star, its crust, covers about one tenth of the total stellar radius and consists of ions, electrons and neutrons. The structure and composition of the crust play an important role in the emission of gravitational waves and the evolution of the neutron star's magnetic field \\citep{brown1998_2,ushomirsky2000,horowitz2009}. Furthermore, studying the crusts of neutron stars allows the investigation of a variety of nuclear reaction processes \\citep{haensel2008,horowitz2008}.  When residing in a low-mass X-ray binary (LMXB), a neutron star strips off the outer gaseous layers of a stellar companion through Roche-lobe overflow, and this matter is accreted onto the neutron star. Transient systems typically exhibit outbursts that last several weeks, after which accretion onto the neutron star ceases and the binary returns to a quiescent state. Usually it remains as such for many years before entering a new active phase. The accretion of matter compresses the neutron star crust, which induces a chain of exothermic electron captures, neutron emissions and density-driven nuclear fusion reactions that take place at several hundreds meters depth \\citep{haensel1990a,haensel2008,gupta07}. The energy released in these reactions locally heats the crust, but is eventually conducted over the entire stellar body and radiated away via neutrinos emitted from the dense core, and via thermal X-ray radiation from the neutron star surface \\citep{brown1998,rutledge2002}. During quiescent episodes, the thermal X-ray radiation emerging from the hot neutron star can be detected with sensitive X-ray satellites.  Prior to the onset of an accretion phase, a neutron star is nearly isothermal and the surface radiation therefore offers a probe of the temperature of the stellar core \\citep{colpi2001}. Shortly after an accretion outburst, however, the thermal X-ray emission tracks the temperature of the heated crust \\citep{ushomirsky2001,rutledge2002}. As the crust cools and restores thermal equilibrium with the core, the thermal X-ray emission decreases and settles at a quiescent base level set by the core temperature \\citep{rutledge2002}. This thermal relaxation depends on the composition of the crust, which determines the conduction of heat towards the surface and the core, and on the details of the nuclear reactions, which set the magnitude and depth of the heat sources \\citep{brown08}.  In the past decade, crust cooling has successfully been monitored for four transiently accreting neutron stars: \\ks\\ \\citep{wijnands2001,cackett2010}, \\mxb\\ \\citep[][]{wijnands2003,cackett2008}, \\xte\\ \\citep[][]{fridriksson2011} and \\exo\\ \\citep[][]{degenaar2010_exo2,diaztrigo2011}. All four belong to a rare sub-group of transient LMXBs, the so-called quasi-persistent sources, which exhibit unusually long accretion episodes extending up to decades. These served as prime targets for crust cooling studies, since the prolonged accretion ensures a substantially heated crust. Comparing the X-ray observations with thermal evolution simulations has provided unique insight into the properties of neutron star crusts \\citep{shternin07,brown08}.  An important conclusion drawn from these studies is that the crust efficiently conducts heat such that the ion lattice must have a highly ordered structure \\citep{shternin07,brown08}. This poses problems for explaining the ignition of thermonuclear superbursts observed from some accreting neutron stars: these require a critical crust temperature that may be difficult to achieve when heat is efficiently conducted \\citep{cooper2005,cumming06}. As a possible solution, it has been hypothesized that accretion deposits more heat inside the neutron star than is accounted for in current models \\citep{brown08,gupta07,horowitz2008}. New observations of neutron star crust cooling can confirm the existence, and constrain the magnitude and location, of such possible additional sources of heat \\citep{brown08}. Unfortunately, X-ray binaries with long accretion phases are very rare: only a handful of sources serve as potential targets for future studies \\citep{wijnands04_quasip}.  \\igr\\ (\\cxo; J1748 hereafter) is a transient LMXB located in the dense core of the globular cluster Terzan 5 (Fig.~\\ref{fig:ds9}). It was initially discovered during \\chan\\ observations of the cluster performed in 2003 and classified as a low-luminosity X-ray source, possibly a dormant X-ray binary \\citep[][]{heinke2006_terzan5}. In 2010 October, the source intensity suddenly increased by orders of magnitude \\citep[][]{bordas2010,pooley2010}. The detection of thermonuclear X-ray bursts \\citep[][]{chenevez2010_terzan5,linares2011} and coherent X-ray pulsations \\citep[][]{strohmayer2010, papitto2010} unambiguously identified J1748 as a neutron star LMXB. Timing analysis of the 11-Hz X-ray pulsations, representing the neutron star's spin period, revealed an orbital period of $\\simeq21$~h and a $\\simeq 0.4$--$1.5 ~\\Msun$ donor star \\citep[][]{papitto2010}.  In 2010 December, $\\sim10$~weeks after its onset, the accretion activity ceased and the binary returned to quiescence \\citep{deeg_wijn2011_2}. A \\chan\\ observation obtained about 50 days later, in 2011 February, demonstrated that the quiescent thermal X-ray emission of J1748 was elevated above the base level measured from archival observations taken in 2003 and 2009, by approximately a factor of 4 \\citep{deeg_wijn2011_2,deeg_wijn2011}. The associated rise in neutron star temperature (a factor of 1.4), can be interpreted as an accretion-heated crust, although two alternative explanations could be invoked \\citep{deeg_wijn2011_2}. First, the neutron star possibly continues to accrete at a very slow rate, causing a variable quiescent flux. Second, the amount of hydrogen and helium left on the neutron star after an accretion phase determine the heat flux that is flowing from the stellar interior towards the surface. For a constant interior temperature, the quiescent thermal emission might therefore differ by a factor of a few due to changes in the envelope composition after an intervening outburst \\citep{brown2002}. Both of these scenarios could also account for the enhanced surface temperature observed in 2011.    \\begin{figure} \\begin{center} \\includegraphics[width=8.0cm]{ds9_small.eps} \\end{center} \\caption[]{{\\chan/ACIS image of Terzan 5 obtained on 2011 April 29--30. J1748 is located in the dense cluster core. }} \\label{fig:ds9} \\end{figure} ",
        "conclusions": "\\label{sec:discuss} We analysed a new \\chan\\ observation of Terzan 5, obtained $\\simeq125$~days after the end of the accretion outburst of J1748. The quiescent X-ray spectrum of the source is soft and can be adequately fit by a neutron star atmosphere model NSATMOS with a temperature of $kT^{\\infty}=92.1\\pm3.1$~eV. This is nearly $10\\%$ lower than measured 75 days earlier in 2011 February ($kT^{\\infty}=100.9\\pm3.3$~eV), yet over $20\\%$ higher than the quiescent base level of the source detected in 2003 and 2009 ($kT^{\\infty}=72.0\\pm3.3$~eV). The thermal X-ray emission and inferred neutron star temperature have thus decreased after the initial enhancement (Fig.~\\ref{fig:lc}).  To explain the elevated thermal emission observed in 2011 in terms of a different envelope composition left after the 2010 accretion phase, the neutron star temperature inferred from our new observation should have been similar to the previous measurement, since the composition of the layer would not change in absence of accretion \\citep[][]{brown2002}. This is not consistent with our results. Continued accretion is thought to reveal itself through stochastic variability in the X-ray emission on both long (months/years) and short (seconds/days) time scales, and/or via the presence of a hard, non-thermal emission component in the X-ray spectrum \\citep{rutledge2001,cackett2010_cenx4,fridriksson2011}. We find none of such features in the \\chan\\ data of J1748, so there are no obvious indications that the neutron star is accreting at a low level.\\footnote{We note that if residual accretion occurs in quiescence, variations in the thermal emission could be due by a changing emission area (a hotspot at the magnetic pole) and may not necessarily imply a changing surface temperature. Given the data statistics, it would not be possible to make this distinction.}  On the other hand, the observed decrease in thermal X-ray emission and neutron star temperature is very reminiscent of the crust cooling observed from the four X-ray binaries with long accretion phases. The data points of J1748 (Fig.~\\ref{fig:lc}) are consistent with an exponential decay with an e-folding time of $\\sim100-200$~days or a power law with a decay index of order $\\sim0.1$. This is similar to the results obtained for the four quasi-persistent LMXBs \\citep{cackett2008,cackett2010,degenaar2010_exo2,fridriksson2011}. We therefore consider this the most likely interpretation of the observations. Thus, we propose that the crust of the neutron star in J1748 was heated during its recent 10-week accretion phase and is currently cooling, supporting previous speculations \\citep{deeg_wijn2011_2,deeg_wijn2011}.   \\subsection{Crust cooling simulations of \\igr}\\label{subsec:sim}  We employed the thermal evolution code of \\citet{brown08}, using all the basic physics ingredients and modeling approach described in that work, to simulate neutron star crust cooling curves for J1748 (Fig.~\\ref{fig:lc}). We assumed that the outburst had a duration of $t_{\\mathrm{ob}}=0.17$~yr and that mass was accreted at a rate of $\\dot{M}=2.0\\times10^{17}$~g~s$^{-1}$, as inferred from X-ray observations \\citep{deeg_wijn2011_2}. Furthermore, the temperature of the neutron star core was fixed at $T=7.0\\times10^{7}$~K, which matches the inferred value of the quiescent base level \\citep{deeg_wijn2011}.  Using standard physics input, the crust temperature is too low and the slope of the cooling curve too flat to match the observations (dotted curve in Fig.~\\ref{fig:lc}). Motivated by recent theoretical conjectures (see Sec.~\\ref{sec:intro}), we also performed simulations with a heat source placed at shallow depth inside the neutron star and varied its strength, $Q_{\\mathrm{extra}}$. We position this extra heat at a column depth of $P/g=4.5\\times10^{11}$~g~cm$^{-2}$, which roughly corresponds to a density $\\rho \\simeq 4\\times10^{8}$~g~cm$^{-3}$ or about 10~m inside the neutron star \\citep{brown08}. This is roughly the depth at which superbursts are thought to ignite. As can be seen in Fig.~\\ref{fig:lc}, adding extra heat raises the temperature and steepens the slope of the crust cooling curve (dashed and solid lines).  For the assumed $\\dot{M}\\simeq2\\times10^{17}~\\mathrm{g~s}^{-1}$ \\citep{deeg_wijn2011}, we find that a heat source of $Q_{\\mathrm{extra}}\\simeq1.0$~MeV per accreted nucleon can match the observational data of J1748 (solid curve in Fig.~\\ref{fig:lc}). This is substantial compared to the total heat release of $\\simeq2.0$~MeV per accreted nucleon that is accounted for by current nuclear heating models \\citep[][]{gupta07,haensel2008}. However, the input value of the mass-accretion rate has a notable effect on the resulting crust cooling curve and hence on the strength of the inferred additional heat source. If we increase $\\dot{M}$ by a factor of two in our model calculations, which is not implausible given the uncertainties in determining this parameter from X-ray observations \\citep{zand07}, the magnitude of the required extra heat decreases to $Q_{\\mathrm{extra}}\\simeq 0.5$~MeV per accreted nucleon. Nevertheless, even if J1748 accreted at the Eddington limit, additional heat is required to explain the observations as cooling of the accretion-heated crust.  As discussed by \\citet{brown08}, heat is primarily conducted in the neutron star crust by electrons, and hence the thermal conductivity is determined by the rate of scatterings between electrons and ions, and between electrons and lattice impurities. The level of impurities in the crust is parametrized by the distribution of ion charges. For the calculations presented in Fig.~\\ref{fig:lc}, this was set to $\\langle Z^2\\rangle - \\langle Z\\rangle^2=5$, which is consistent with the values inferred from modeling the crust cooling curves of two other neutron stars \\citep{brown08} and molecular dynamics simulations \\citep{horowitz2009_cond}. We explored simulations with different levels of impurities, but found that this does not affect the resulting crust cooling curve of J1748. This is because, unlike systems with accretion episodes that last for years and heat the entire crust, for short outbursts the heating primarily occurs at low densities, where the heat conduction is controlled by electron-ion scattering.  We calculated crust cooling curves varying the depth of the extra heat source. We found that the exact depth primarily affects the cooling rate at early times, within a few tens of days after the end of the accretion outburst. J1748 could not be observed with \\chan\\ on such a time scale, because the source was too close to the Sun at that time. Thus, the location of the extra heat source does not affect the crust cooling curve at the time probed by the observations of J1748. This also implies that our \\chan\\ observations cannot firmly constrain the depth of the additional heat release, although we can obtain some limits.  Assuming that the crust cooling curve of J1748 indeed follows a power law decay with an index of $-0.1$ (see above), we estimate that there is an inward-directed heat flux in the crust of $F\\simeq3\\times10^{21}~\\mathrm{erg~cm}^{-2}~\\mathrm{s}^{-1}$ \\citep[][eq. 12]{brown08}. At any given depth inside the neutron star, it takes a certain characteristic time for heat to diffuse to the surface. Therefore, at a given time there will be a certain depth at which the layer above has already thermally relaxed, while deeper layers are still hot and have not yet started to cool \\citep{brown08}. The first \\chan\\ observation of J1748 was obtained $\\simeq50$~days after the end of the accretion outburst: this is the characteristic thermal time scale for layers located at a column depth of roughly $P/g\\simeq10^{14}~\\mathrm{g~cm}^{-2}$, which corresponds to a matter density of $\\rho \\simeq 3\\times10^{10}~\\mathrm{g~cm}^{-3}$ or a depth of $\\simeq150$~m \\citep{brown08}. Since the heat flux is directed inwards, there must thus be a source of heat exterior to the depth corresponding to the thermal time scale of $50$~days. This implies that there is a heat deposit at less than $\\simeq150$~m inside the neutron star \\citep{brown08}. This is much shallower than the depth at which most crust nuclear reactions are thought to occur \\citep{gupta07,haensel2008,horowitz2008}.   \\begin{table} \\caption{Results from fitting the \\chan\\ spectral data.} \\begin{threeparttable} \\begin{tabular}{c c c c c c c} \\hline \\hline Date & $kT^{\\infty}$ & $F_{\\mathrm{unabs}}$ & $L_{\\mathrm{bol}}$\\\\ \\hline 2003-07-13/14 & \\multirow{2}{*}{$72.0\\pm3.3$} & \\multirow{2}{*}{$1.64\\pm0.30$} & \\multirow{2}{*}{$0.59\\pm0.11$} \\\\ 2009-07-14/15 &  &  \\\\ 2011-02-17 & $100.9\\pm3.3$ & $6.33\\pm0.83$ & $2.29\\pm0.30$ \\\\ 2011-04-29/30 & $92.1\\pm3.1$ & $4.39\\pm0.60$ & $1.59\\pm0.21$ \\\\ \\hline \\end{tabular} \\label{tab:spec} \\begin{tablenotes} \\item[]Note. $kT^{\\infty}$ is given in eV, $F_{\\mathrm{unabs}}$ in $10^{-13}~\\flux$ (0.5--10 keV) and $L_{\\mathrm{bol}}$ in $10^{33}~\\lum$ (0.01--100 keV). \\end{tablenotes} \\end{threeparttable} \\end{table}   \\subsection{Repercussions}\\label{subsec:reper} The X-ray observations of J1748 obtained after its recent 10-week accretion outburst can be explained as cooling of the accretion-heated neutron star crust. However, regardless of the exact value of $\\dot{M}$, this requires a substantial source of heat located at shallow depth inside the neutron star that is not accounted for by existing nuclear physics calculations \\citep{gupta07,haensel2008,horowitz2008}. A shallow heat source has also been proven necessary to explain the crust cooling curve of at least one of the quasi-persistent LMXBs \\citep[\\mxb;][]{brown08}. The presence of substantial heat sources in the outer crustal layers is important for the ignition of superbursts: the extra heat release raises the crust temperature and can thus compensate for the effect of a high thermal conductivity \\citep{gupta07,brown08}.  Our work on J1748 suggests that the crust of a neutron star can become substantially heated during an accretion episode of only a few weeks, and that it is technically feasible to detect the subsequent crust cooling with existing, sensitive X-ray instruments. LMXBs in which the neutron star accretes for weeks or months are much more common than those that undergo extended accretion episodes of years. Future crust cooling studies can therefore draw from a larger sample of transiently accreting neutron stars. The most promising targets are neutron stars that have relatively low core temperatures, but exhibit bright accretion episodes so that a substantial amount of matter is consumed and a considerable temperature gradient can develop in the crust \\citep{brown1998}. Transient Be/X-ray binaries, which harbor strongly magnetized neutron stars feeding of the matter expelled by their massive Be companions, can potentially also be included in future studies \\citep{brown1998}.  Studying a larger sample of crust cooling curves allows to investigate whether the presence of strong, shallow heat sources is a common feature amongst transiently accreting neutron stars. In particular, observations performed within $\\sim1$~month after the end of an accretion outburst can further constrain the magnitude and depth of such heat sources \\citep{brown08}. The next challenge is then to identify the nature of this extra heat release.  Our crust cooling simulations suggest that \\igr\\ continues to cool within the next year. New \\chan\\ observations are planned to confirm this and to put further constraints on the shape of the crust cooling curve.   \\begin{figure} \\begin{center} \\includegraphics[width=8.0cm]{crust_curves.eps} \\end{center} \\caption[]{{ Evolution of the neutron star temperature of \\igr, as inferred from analysis of \\chan\\ spectral data. The grey shaded area marks the quiescent base level observed prior to the accretion activity of 2010 October--December. The three curves indicate neutron star thermal evolution simulations for different amounts of heat production in shallow crustal layers. }} \\label{fig:lc} \\end{figure}   ~\\\\ \\noindent {\\bf Acknowledgements.} The authors are grateful to Harvey Tananbaum and the \\chan\\ science team for making this observation possible. RW acknowledges support from a European Research Council (ERC) starting grant. Support for EFB was provided by NASA through \\chan\\ Award Number TM0-11004X. This work was supported by the Netherlands Research School for Astronomy (NOVA). ND and EFB acknowledge the hospitality of the International Space Science Institute (ISSI) in Bern, Switzerland, where part of this work was carried out.  \\vspace{-0.5cm}"
    },
    "1107/1107.5584_arXiv.txt": {
        "abstract": "We apply Support Vector Machines -- a machine learning algorithm --  to the task of classifying structures in the Interstellar Medium. As a case study, we present a position-position velocity data cube of $^{12}$CO J=3--2 emission towards \\snr{}, a supernova remnant that lies behind the M17 molecular cloud. Despite the fact that these two objects partially overlap in position-position-velocity space, the two structures can easily be distinguished by eye based on their distinct morphologies. The Support Vector Machine algorithm is able to infer these morphological distinctions, and associate individual pixels with each object at $>$90\\% accuracy. This case study suggests that similar techniques may be applicable to classifying other structures in the ISM -- a task that has thus far proven difficult to automate. ",
        "introduction": "\\label{sec:intro}  Classifying interesting objects in a dataset is an essential early step in most analysis tasks. For many objects in astronomy (stars, galaxies, solar system objects, extragalactic supernovae), algorithms can identify and characterize these objects automatically \\citep{Irwin85, Bertin96, Naylor98}. Structures in the Interstellar Medium (ISM), however, have proven difficult to classify in this way. These objects -- which include, e.g., molecular clouds, infrared dark clouds, bubbles, jets, radiation-shaped pillars, and filaments -- are morphologically complex and heterogeneous. The essential properties of these structures are hard to encode.  To take advantage of increasingly wide-area surveys, previous researchers have mainly relied on manual identification of features in the ISM \\citep{Churchwell06, Helfand06, Curtis10, Arce10}. There are several drawbacks to this approach: it is time consuming, non-repeatable and, when identifying complex structures, affected by difficult-to-quantify selection effects related to how specific people perceive an image. Despite recent advances in the study of specific objects (e.g. Infrared Dark Clouds, \\citealt{Peretto09};  filaments, \\citealt{Menshchikov10}), automated feature identification in the ISM is an still an open problem.  Machine learning algorithms are designed to infer patterns in data which are otherwise difficult to define explicitly. They are grouped into two classes: supervised (in which the algorithm is ``trained'' to recognize a pattern via a set of training examples) and unsupervised (in which the algorithm identifies groupings within a dataset \\textit{a priori}).  These algorithms can mechanize the process of object identification, and hence address many of the shortcomings of manual  classification -- they scale easily to other similar data, and their results are repeatable. A potential drawback of these methods is that, because the classification is not guided by a physical model, they are susceptible to over-fitting and to inheriting selection biases within the training data. However, because the machine approach extends easily to other similar data, these biases are more readily characterized via classification of test data sets.  In this paper, we explore whether machine learning algorithms can be used to catalog structures in the ISM. We present as a case study a spectral line data cube of $^{12}$CO emission towards the M17 star forming region. Emission from this cloud overlaps with \\snr{}, a supernova remnant situated behind the cloud. We use the Support Vector Machine (SVM, \\citealt{Vapnik95}) algorithm to identify the supernova remnant and, on a pixel-by-pixel basis, classify the original data cube. We are able to use this classification to derive the mass and momentum of the supernova remnant, which would not otherwise be possible with these data. ",
        "conclusions": "We have presented a case study of a supervised learning task applied to classifying structures in the ISM. The M17 molecular cloud overlaps shocked CO emission from \\snr{} but, because of the supernova's distinct morphology in position-position-velocity space, the two objects are readily distinguishable by eye. The SVM classification algorithm is able to learn these morphological differences using a representative sample of manually-classified pixels. We emphasize several important characteristics of this approach:  \\begin{enumerate} \\item Machine-based classification of ISM structures permits a pixel-level classification of datasets. This level of refinement is often prohibitively cumbersome via manual identification. \\item By using an independent set of manually-classified validation data, we can characterize the quality of this classification. We can further use this information to refine and improve the algorithm's performance. \\item Extracting information about the moments of the intensity distribution in our data produced the most effective classification. Other information (the weights in a principal component analysis, spatial derivatives) was less successful. \\item Only a very small fraction of the data ($\\sim0.1\\%$) needs to be categorized to train the algorithm. An efficient approach for future work may be to evaluate the classifier's performance as the training set is assembled, to better understand when the training set is large (and representative) enough. \\end{enumerate}  This case study suggests that automated algorithms are capable of identifying complex structures seen in the ISM. Such an approach may be useful in analyses of current and future surveys of the Milky Way's ISM, particularly when identifying morphologically distinct structures like bubbles, pillars, and filaments.  We are grateful to C. Brogan for sharing her 20\\,cm data of \\snr{}, and to K. Binsted for discussions. We also thank the anonymous referee, whose careful reading and comments improved the clarity of this paper.  This material is based upon work supported by the National Science Foundation under Grant No. AST-0908159.  \\pagebreak"
    },
    "1107/1107.1673.txt": {
        "abstract": "\\rxja{} and \\rxjn{} are the only young isolated radio-quiet neutron stars (NSs) for which trigonometric parallaxes were measured. Due to detection of their thermal emission in X-rays they are important to study NS cooling and to probe theoretical cooling models. Hence, a precise determination of their age is essential.\\\\ Recently, new parallax measurements of \\rxja{} and \\rxjn{} were obtained. Considering that NSs may originate from binary systems that got disrupted due to an asymmetric supernova, we attempt to identify runaway stars which may have been former companions to the NS progenitors. Such an identification would strongly support a particular birth scenario with time and place. \\\\ We trace back each NS, runaway star and the centres of possible birth associations (assuming that most NSs are ejected directly from their parent association) to find close encounters. The kinematic age is then given by the time since the encounter. We use Monte Carlo simulations to account for observational uncertainties and evaluate the outcome statistically.\\\\ Using the most recent parallax measurement of $\\unit[8{.}16\\pm0{.}80]{mas}$ for \\rxja{} by \\citeauthor{2010ApJ...724..669W}, we find that it originated in the Upper Scorpius association $\\unit[0{.}46\\pm0{.}05]{Myr}$ ago. This kinematic age is slightly larger than the value we reported earlier ($\\unit[0{.}3]{Myr}$) using the old parallax value of $\\unit[5{.}6\\pm0{.}6]{mas}$ by \\citeauthor{2003conf...Kaplan}. Our result is strongly supported by its current radial velocity that we predict to be $\\unit[6^{+19}_{-20}]{km/s}$. This implies an inclination angle to the line-of sight of $\\unit[88\\pm6]{deg}$ consistent with estimates by \\citeauthor{2001A&A...380..221V} from the bow shock. No suitable runaway star was found to be a potential former companion of \\rxja{}.\\\\ Making use of a recent parallax measurement for \\rxjn{} of $\\unit[3{.}6\\pm1{.}6]{mas}$ by \\citeauthor{ThomasPhD}, we find that this NS was possibly born in Trumpler 10 $\\unit[0{.}85\\pm0{.}15]{Myr}$ ago. This kinematic age is somewhat larger than the one obtained using the old parallax value of $\\unit[2{.}77\\pm1{.}29]{mas}$ by \\citeauthor{2007ApJ...660.1428K} ($\\unit[0{.}5]{Myr}$). We suggest the B0 runaway supergiant HIP 43158 as a candidate for a former companion of the progenitor star. Then, the current distance of \\rxjn{} to the Sun should be $\\unit[286^{+27}_{-23}]{pc}$, in agreement with recent measurements. We then expect the radial velocity of \\rxjn{} to be $\\unit[-76^{+34}_{-17}]{km/s}$. ",
        "introduction": "\\label{sec:intro}  Neutron stars (NSs) show large proper motions which, with known distances, indicate high space velocities \\citep[e.g.][]{1994Natur.369..127L,1997MNRAS.289..592L,1997MNRAS.291..569H,1998ApJ...505..315C,2002ApJ...568..289A,2005MNRAS.360..974H}. Some NSs even show velocities of the order of $\\approx\\unit[1000]{km/s}$ (e.g. PSR B1508+55, \\citealt{2005ApJ...630L..61C}; PSR B2223+65, \\citealt{1993MNRAS.261..113H}; \\citealt{1993ApJ...411..674T}; RX J0822-4300, \\citealt{2006A&A...457L..33H}, \\citealt{2007ApJ...670..635W}). Those high velocities may be the result of asymmetric supernova (SN) explosions assigning the new-born NS a kick velocity for that a number of mechanisms have been suggested \\citep[e.g.][]{1996PhRvL..76..352B,1996A&A...306..167J,2005ASPC..332..363J,2006ApJ...639.1007W,2009arXiv0906.2802K}. Another possibility is that the high-velocity NSs are the remnants of (symmetric\\footnote{Although numerous three-dimensional simulations showed that SNe are most likely asymmetric \\citep[e.g.][]{2010ApJ...725L.106W,2010ApJ...714.1371H,2008PhRvD..78f4056D,2008AIPC..983..369J,2006A&A...457..963S,2007Natur.445...58B}.}) SN explosions of the so-called hyper-velocity runaway stars \\citep{2007A&A...470L...9G,2008MNRAS.385..929G,2009MNRAS.395L..85G} which were ejected due to dynamical three- or four-body encounters either from the Galactic Centre \\citep{1988Natur.331..687H} or from massive star clusters in the Galactic disk.\\\\ About $\\unit[30]{\\%}$ of young stars show different velocity properties than normal Population I stars \\citep[][thereafter T11, a catalogue of young runaway stars]{1991AJ....102..333S,1993ASPC...35..207B,2011MNRAS.410..190T}. Two scenarios are accepted to produce those so-called ``runaway stars'' \\citep{1961BAN....15..265B}. The binary-SN scenario \\citep{1961BAN....15..265B} is related to the formation of the high velocity NSs (but note that high-velocity NSs may also be the result of an asymmetric SN explosion of a single massive star or a SN of a massive runaway star, see above): The runaway and NS are the products of a SN within a binary system. The velocity of the former secondary may be as large as its original orbital velocity \\citep{1998A&A...330.1047T}. Runaway stars produced in this scenario should share typical properties such as a high rotational velocity $v\\sin i$ and an enhanced helium abundance owing to momentum and mass transfer during binary evolution \\citep{1993ASPC...35..207B}. We refer to such runaway stars as BSS runaway stars. The second scenario is the dynamical ejection due to gravitational interactions between massive stars in dense clusters \\citep{1967BOTT....4...86P}.\\\\ Several studies have been made to investigate the origin of runaway stars \\citep[e.g.][]{1961BAN....15..265B,1986ApJS...61..419G,2001A&A...365...49H,2005A&A...437..247D,2008A&A...489..105S,2008A&A...490.1071G} but only a few regarding the origin of fast-moving NSs \\citep[e.g.][thereafter T10]{2001A&A...365...49H,2004ApJ...610..402V,2010MNRAS.402.2369T}.\\\\ In this paper we will re-investigate the origin of two isolated radio-quiet X-ray emitting NSs, namely \\rxja{} and \\rxjn{}. So far, only seven such sources have been identified for which they were named ``The Magnificent Seven'' (M7) (\\citealt{2001ASPC..234..225T}; for recent reviews see \\citealt{2007Ap&SS.308..181H,2008AIPC..968..129K}). They are bright X-ray sources associated with faint blue optical counterparts arising from a hot cooling surface. In only the two cases studied in this paper a trigonometric parallax was obtained. The M7, and especially \\rxja{} and \\rxjn{}, are very important NSs since brightness, parallax and temperature yield the size of the emitting area and hence, their (model dependent) radii. From spectra, one can in principle determine their masses and atmosphere composition, which eventually may lead to constraints on the equation of state.  With known luminosity and age, cooling curves can be verified. As the characteristic age represents only a rough estimate of the true age \\citep[e.g.][]{1988MNRAS.234P..57B,2000Natur.406..158G,2002ApJ...567L.141M}, the kinematic age is very important to have a better estimate of the true age. Also, the characteristic age is significantly influenced by pulsar winds \\citep[e.g.][]{2003A&A...409..641W} and possibly by emission of gravitational waves \\citep[e.g.][]{2008CQGra..25w5011W}.\\\\ Under the assumption that most NSs form in associations or clusters of massive stars, the identification of the birth association of a NS by its flight path is possible if the uncertainties in the distance are moderate. Our assumption is justified since we observe associations and clusters of massive stars (also, the dispersion time scale of a massive star cluster is much longer than the lifetime of its most massive member stars), i.e. only a small fraction of their massive member stars is ejected from their parent cluster due to gravitational interactions before they end their lives in SNe (about 20-30 per cent of O and early B type stars are located outside of clusters, \\citealt[e.g.][]{1998AJ....115..821M,2004ApJS..151..103M}, i.e. $\\gtrsim$70 per cent of a cluster's member stars remain in the cluster). There are certainly NSs that form outside their parent cluster as there are massive runaway stars that will later explode in a SN event and become NSs. For the latter, the identification of their parent associations or birth sites is hardly possible. Also, it is possible, that a massive binary was ejected from its parent cluster. If the primary experiences a SN, the secondary might be ejected from the system and will later explode in a SN. For NSs whose progenitors experienced this so-called two-step-ejection scenario \\citep{2010MNRAS.404.1564P}, it is not possible to identify their formation sites. However, this scenario applies only to 1-4 per cent of O stars \\citep{2010MNRAS.404.1564P}. NSs formed from runaway stars may show a higher space velocity than those formed in their parent cluster since their velocity vector is a superposition of the runaway star's velocity and the kick velocity the NS receives at birth.\\\\ When searching for the parent association of a NS (under the above assumptions), the result is often not unique (T10). For that reason, it is desirable to find a second indicator for the location and time of the past SN, e.g. a possible former companion star that is now a runaway star.\\\\ Here, we note that not for every NS a (identifiable) runaway star must exist. The NS progenitor could have been a single star (possibly a runaway star, see above); a former massive companion could have already undergone a SN and is now a NS; or the former companion that is now a runaway star was not yet identified as such because, e.g., its absolute velocity is low \\citep[e.g.][]{1998A&A...330.1047T} or the direction of its velocity vector is not significantly different from those of its neighbouring stars.\\\\  After investigating possible parent associations for \\rxja{} and \\rxjn{} in \\autorefs{subsec:1856parentassoc} and \\ref{subsec:0720parentassoc}, respectively, applying the procedure described in \\autoref{sec:procedure} and using most recent parallax measurements, we attempt to identify the possible former companion for \\rxja{} and \\rxjn{} in \\autorefs{subsec:1856comp} and \\ref{subsec:0720comp}, respectively. We summarise our results and draw our conclusions in \\autoref{sec:summary}.  %__________________________________________________________________ ",
        "conclusions": "\\label{sec:summary}  We analysed the origin of two members of the M7, \\rxja{} and \\rxjn{} using most recent parallax measurements. \\\\ Under the assumption that most NSs are born in and ejected from their parent association or cluster (see \\autoref{sec:intro}), we confirm that \\rxja{} most probably originated from the US association as suggested previously (\\citealt{2002ApJ...576L.145W}, T10). We find that the radial velocity of \\rxja{} is $\\unit[6^{+19}_{-20}]{km/s}$ implying an inclination angle to the line-of-sight of $\\unit[88\\pm6]{deg}$ that is consistent with the inclination of the bow shock that \\rxja{} creates in the ISM owing to its motion \\citep{2001A&A...380..221V}. This consistency implies that the current distance of \\rxja{} is $\\approx\\unit[120]{pc}$ as obtained by \\citet{2002ApJ...576L.145W} and \\citet{2010ApJ...724..669W} rather than $\\unit[160-180]{pc}$ \\citep{2003conf...Kaplan,2007Ap&SS.308..191V}. This smaller distance is also in good agreement with \\rxja{} being in front of the Corona Australis star forming region ($\\approx\\unit[130]{pc}$, \\citealt{2008hsf2.book..735N}); if it was behind, optical detection would probably have failed. Moreover, this is in accordance with the absence of OI from the ISM in the Chandra HRC/LETG spectra \\citep{2001A&A...379L..35B,2003A&A...399.1109B}. The space velocity of \\rxja{} would then be $\\approx\\unit[195]{km/s}$. The derived kinematic age of \\rxja{} of $\\unit[0{.}46\\pm0{.}05]{Myr}$ agrees well with theoretical cooling models (cf. T10, Figure 12 therein). \\\\ The mass of the progenitor star would have been $\\approx\\unit[40-60]{M_\\odot}$. Considering that the NS may have originated from a binary system that got disrupted in an asymmetric SN explosion, we searched for a possible former companion that should then be a runaway star. No former companion candidate was identified. A reason could be that, if there was a former companion, either, it may not be included in the runaway star catalogue of T11 either because it was no Hipparcos source or its velocity vector is not significantly different from those of its neighbouring stars (note, that T11 also investigated the direction of motion of young stars compared to neighbouring stars, regardless whether the absolute velocity is large or small); or a former companion might also have experienced a SN after ejection (then, the mass ratio of the binary would have been close to unity). Alternatively, the progenitor star of \\rxja{} could have been a single star and the NS gained its high velocity in an asymmetric SN explosion. However, single stars of masses $\\gtrsim\\unit[20-25]{M_\\odot}$ are generally believed to produce black holes rather than NSs. \\citet{2001ApJ...560L..83M}, \\citet{2002A&A...390..299B} and \\citet{2009SSRv..143..437F} suggested that about 20 nearby (in Sco-Cen) SNe created and re-heated the Local Bubble. \\rxja{} was probably formed in one of them about half a million years ago in US. For that reason and since the small radial velocity of \\rxja{} that we found is strongly supported by the bow shock measurements of \\citet{2001A&A...380..221V}, we argue that \\rxja{} is not a remnant of a runaway star that had already left its parent cluster when it experienced a SN.\\\\  For \\rxjn{} there is no unique result on the parent association due to the large parallax uncertainties (as long as runaway stars are not concerned). It could have been formed either in one of the young local associations \\citep{2008A&A...480..735F} with TWA being the best candidate birth place (see also T10) or in a more distant association/cluster such as Tr 10 or Col 140. In the first case, the distance of the SN to the Sun is very small ($\\approx\\unit[30-50]{pc}$) and it is expected to find SN-produced radionuclides on Earth \\citep{1996ApJ...470.1227E}. A small but insignificant signal of $^{60}$Fe was found for $\\unit[0-1]{Myr}$ in the past \\citep{2004PhRvL..93q1103K}. \\\\ The identification of another indicator is thus needed. We find three runaway stars for which not only a positional coincidence with \\rxjn{} is possible but also an enhancement of $^{26}$Al emission was measured at the location of the potential SN \\citep{2010A&A...522A..51D}. Regarding the current distance of \\rxjn{} derived for each case, we suggest that \\rxjn{} may have been born $\\unit[0{.}85\\pm0{.}15]{Myr}$ ago as a former member of Tr 10 with HIP 43158 being the possible former companion. The mass of the progenitor star would have been $\\approx\\unit[8-14]{M_\\odot}$. The current distance of \\rxjn{} would then be $\\approx\\unit[285]{pc}$, in good agreement with estimates from $n_H$ measurements \\citep[$250\\pm25\\,\\mathrm{pc}$,][]{2007Ap&SS.308..171P} and the most recent HST parallax (distance $\\unit[280^{+210}_{-85}]{pc}$, \\citealt{ThomasPhD}). The present radial velocity of \\rxjn{} would be $\\unit[-76^{+34}_{-17}]{km/s}$ leading to a space velocity of $\\approx\\unit[160]{km/s}$.\\\\ However, we emphasise that further observation of the former companion candidate is needed to confirm its status. There are also runaway stars without radial velocity measurements that are potential former companion candidates for \\rxjn{}. For them, a measurement of the radial velocity is crucial.\\\\ \\rxjn{} is probably older than \\rxja{}, thus should be cooler than \\rxja{}. However, the effective temperature of \\rxjn{} is $kT\\approx\\unit[85-95]{eV}$ \\citep{2009A&A...498..811H} while for \\rxja{} it is $kT\\approx\\unit[60]{eV}$ \\citep{2003A&A...399.1109B}. Hence, \\rxja{} is cooler although it is probably younger. This might be due to different birth parameters or different cooling. Moreover, effective temperatures may be influenced by hot spots, hence do not reflect the real surface temperature. For \\rxjn{}, this is indeed the case \\citep{2006A&A...451L..17H,2009A&A...498..811H} and it appears hotter."
    },
    "1107/1107.1855_arXiv.txt": {
        "abstract": "We have produced the next generation of quasar spectral energy distributions (SEDs),  essentially updating the work of Elvis et al.\\ (1994) by using high-quality data obtained with several space and ground-based telescopes, including NASA's Great Observatories. We present an atlas of SEDs of 85 optically bright, non-blazar quasars over the electromagnetic spectrum from radio to X-rays.  The heterogeneous sample includes 27 radio-quiet and 58 radio-loud quasars. Most objects have quasi-simultaneous ultraviolet-optical spectroscopic data, supplemented with some far-ultraviolet spectra, and more than half also have \\spitzer\\ mid-infrared IRS spectra.  The X-ray spectral parameters are collected from the literature where available.  The radio, far-infrared, and near-infrared photometric data are also obtained from either the literature or new observations.  We construct composite spectral energy distributions for radio-loud and radio-quiet objects and compare these to those of Elvis et al., finding that ours have similar overall shapes, but our improved spectral resolution reveals more detailed features, especially in the mid and near-infrared. ",
        "introduction": "The supermassive black holes powering quasars (or active galactic nuclei (AGN)) do not themselves shine.  It is the heated material surrounding the black holes that emits the radiation signatures of quasars.  These signatures include broad emission lines characteristic of high-velocity gas moving at thousands of kilometers per second, and extremely high continuum luminosity in excess of that of entire galaxies.  Continuum emission is seen in all parts of the electromagnetic spectrum from the highest energies (gamma rays and X-rays) to the lowest (radio waves).  The power emitted is similarly high, within an order of magnitude or so, over much of this range, although there can be significant variation from quasar to quasar.  There are no single states of matter or single processes capable of reproducing the spectral energy distribution (SED) of a quasar. A combination of both thermal and non-thermal processes has been invoked to explain, at different parts of the SED, the emission from gas in a variety of states, at a variety of temperatures, at a variety of distances, and experiencing a variety of environments. Quasars seem to have many different components that are expressed in different parts of the electromagnetic spectrum.     So what are the components of the quasar?  The central supermassive black hole is the ultimate engine, allowing the liberation of gravitational potential energy.  The primary source of electromagnetic emission is likely an accretion disk formed of hot gas spiraling into the black hole and shining in the optical through ultraviolet (UV).  A hot atmosphere upscatters the disk photons to X-ray energies. A jet shooting out from the inner accretion disk emits synchrotron radiation, dominating radio emission and sometimes higher energies. An obscuring torus of relatively cool gas and dust, heated by photons from the accretion disk, thermally radiates in the near-infrared (NIR) and mid-infrared (MIR).  The far-infrared (FIR) part of the SED comes from cooler dust, perhaps distributed throughout the host galaxy, that may be heated by stars rather than the quasar itself.  Other regions in and among these continuum-emitting parts are responsible for the prominent emission lines present in quasar spectra. The details of all these parts, mechanisms, and their relationships are not yet completely understood, because multi-wavelength data to further our understanding have been difficult to gather.  Multiwavelength astronomy is challenging.  No single telescope can observe at all wavelengths.  Many parts of the electromagnetic spectrum cannot penetrate the Earth's atmosphere and require space-based observations.  Quasars are variable and this means that simultaneous or nearly simultaneous observations are desirable, at least in some parts of the spectrum.  Different technologies have different sensitivity levels and what may be an easily observed target at one wavelength may be difficult to detect at another.  There were a number of pioneering works on quasar SEDs  in the 1980s \\citep[e.g.,][]{EdeMal86,Ward87,Kriss88,Sande89,SunMal89}.  In the 1990s, \\citet[][hereafter E94]{Elvis94} established the first large, high-quality atlas of quasar SEDs. The timing of their work was predicated on the launch of several space-based telescopes, notably IRAS in the mid to far infrared, IUE in the UV, and Einstein in the X-rays, that for the first time provided observations of large numbers of quasars in these wavebands.  They established the differences between the SEDs of radio-loud (RL) and radio-quiet (RQ) quasars.  They also characterized the variance of quasar SEDs, which is rather substantial, and explored the problem of bolometric corrections.  Since the work of E94, there have been several significant investigations of quasar SEDs \\citep[e.g.,][]{Kuras03,RisElv04}.  \\citet{Richa06} is the largest that covers the entire electromagnetic spectrum, including data from the Sloan Digital Sky Survey and \\spitzer, supplemented by near-IR, GALEX ultraviolet, VLA radio, and ROSAT X-ray data, where available. One of their key findings was again the wide range of SED shapes, and how assuming a mean SED can potentially lead to errors in bolometric luminosities as high as 50\\%.  Other investigations of SEDs have focused on subclasses, like broad absorption line (BAL) quasars \\citep[e.g.,][]{Galla07}, 2MASS red quasars \\citep{Kuras09}, hard X-ray selected quasars \\citep{Polle00, Polle07}, or on individual objects \\citep[e.g.,][]{Zheng01}.  More detailed SED work has been done on limited portions of the entire electromagnetic spectrum, such as the optical-X-ray \\citep[e.g.,][]{Laor97}, the FIR to optical \\citep[e.g.,][]{Netze07}, or optical/UV to X-ray \\citep{Grupe10}.   There has also been work looking for relationships among more detailed spectral features, like emission lines and SEDs.  \\citet{Wilke99}, examining 41 quasars with SED information, for instance, found a variety of Baldwin effects (Baldwin 1977), anticorrelations between emission-line equivalent width and UV luminosity, as well as some correlations between properties of Fe II and C IV emission lines and the ratio of the optical to X-ray luminosity.  \\citet{Schwe06} and \\citet{Netze07} studied Palomar-Green quasars with FIR to optical data and confirmed that most FIR radiation is due to star-forming activity.  However, \\citet{Netze07} also argue, based on a correlation between $L(5100\\mbox{\\AA})$ and $L(60\\micron)$, an alternative view that a large fraction of FIR radiation could result from direct AGN heating.  Ideally what one would like in studying quasars is a complete inventory over all time and all directions of all photons emitted. The best we can do now is to obtain spectrophotometric snapshots, close in time, of some spectral regions, supplemented by photometry in other accessible wavebands, from our particular line of sight toward a quasar. The technology has improved since the 1990s, with spacecraft such as HST, Chandra and XMM, and Spitzer replacing IUE, Einstein and ROSAT, and IRAS, respectively.  This atlas is meant to update the last decade's work with a modern set of quasar SEDs using the next generation of telescopes and instruments. Whenever possible we have striven to use high-quality spectrophotometry in addition to photometry, which will enable investigations like those of Wilkes et al. (1999).  One hope is that spectral features will be found that correlate with SEDs and allow their shapes to be determined without the need for complete multi-wavelength observations. Any correlations found should provide deeper insight into quasars. We plan additional papers based on this data set performing various types of investigations, as well as addressing both observational and theoretical aspects of bolometric corrections.  In the next section (\\textsection\\ref{sec:sample}), we describe our sample, which is composed of three subsamples that have excellent quasi-simultaneous optical through UV or far-UV (FUV) spectrophotometry serving as a starting point for SED construction.  We then describe the sample properties.  Subsequent sections describe the data (\\S\\ref{sec:data}), starting with about 100~MHz radio and moving to higher energies to about 10\\,\\kev\\ X-ray.  We then discuss corrections to the SEDs, such as Galactic dereddening and host galaxy removal (\\S\\ref{sec:correction}). In \\S\\ref{sec:sed}, we present the SEDs for our quasars individually, then as composites for radio-loud and radio-quiet subsamples, which are known to differ significantly. We then discuss the properties of the SEDs, and compare our composites with those of E94.  We finish this paper with a summary of these results, future plans, and some concluding remarks (\\S\\ref{sec:summary}). ",
        "conclusions": "}  We try to keep all the original data in building the SEDs of individual objects.  The only change is the re-sampling of the UV-optical spectra to 10\\,\\AA\\ resolution by rebinning.  Although we lose some useful information (e.g. resolving narrow emission lines), this does not affect the SED work at all.  In constructing the composite SEDs, we applied rebinning again mostly for spectroscopic data.  We did not apply any smoothing or interpolation in regions with real data, which could introduce systematic biases. The features in our median SEDs are real.   At the edges of some wavebands, the number of objects with data drops sharply (Fig.~\\ref{fg:msed}, middle), and our method of using median to build the composite SEDs can help to some extent in the small number statistics. We also visually check to ensure that the SEDs are reasonably smooth in those regions.  \\subsubsection{RL vs. RQ\\label{sec:lowzrl}}  Figure~\\ref{fg:lowzrl} overplots our median SEDs of RL and RQ samples, normalized at 4215\\,\\AA\\ during the construction.  The SEDs from FIR to UV are very similar for RL and RQ, especially in the UV-optical region.  This is only true for the UV-optical continuum, because emission lines, such as \\feii\\ and \\oiii, are known to strongly correlate with radio loudness in the Eigenvector I relationship \\citep{BG92}.  We will investigate the relationships between SEDs and emission line properties in a future paper.  The biggest difference between RL and RQ median SEDs is in the radio, where luminosity could differ by 3 orders of magnitude. There is also an obvious difference in the X-ray, where the RL objects are more X-ray luminous.  This correlation between radio and X-ray luminosity has been reported in previous studies \\citep[e.g.,][]{Brink00,Polle07}.  During the construction of the median SEDs, we have also defined 6 bins in the radio for the RQ sample (Table~\\ref{tb:bins}), however there are not enough objects ($\\geq 8$) in 4 of the bins.  Therefore, there are only two points good enough to be included in the final radio SED of RQ sample.  Given the big difference in the number of objects involved in these two bins (Fig.~\\ref{fg:msed}), the apparent steep RQ spectral index, if defined using the two points, may not be reliable.  Our RL sample has more objects with higher redshifts than the RQ sample (Fig.~\\ref{fg:hist}).  We therefore build another median SED of RL objects only with redshift less than 0.5 (21 objects), similar to those in the RQ sample.  Comparing this and the SED of the entire RL sample (Fig.~\\ref{fg:lowzrl}), we find that the only notable difference is in the radio where the low-z subsample shows a less luminous radio SED, but radio spectral index seems similar. The difference is probably real because of sample properties --- high-z RL objects are more luminous in radio, but the differences between the RL and RQ samples are still much more prominent.   \\subsubsection{Comparison with E94 Mean SEDs}  \\citet{Elvis94} use 47 objects to build the mean SEDs (MSED94) for RL and RQ objects.  There are 11 objects in common with our sample, including 6 RL and 5 RQ objects.  We compare our median SEDs of RL and RQ quasars (Fig.~\\ref{fg:msed}) with MSED94.  The overall shape of the SEDs over the entire available frequency range is similar, but there are more detailed features in our new SEDs.  We specifically keep all the emission features in the UV-optical region because they are real spectral features and our data quality allows us to keep them.  The underlying continuum shapes in this region look similar to those of MSED94.  We note that our UV-optical SEDs extend to shorter wavelength beyond 1000\\,\\AA, and start to turn over, indicating the peak of the ``big blue bump'' \\citep[e.g.,][]{Zheng98,Shang05}.  This is especially obvious in our RL median SED where we have more higher redshift objects.  In the MIR, the broad silicate emission features around 10 and 18\\micron\\ are prominent in the SEDs.  These could not otherwise be reproduced without the unprecedented spectral data from \\spitzer\\ IRS.  To the shorter wavelength of these features, a well-defined power-law rises up to about 4\\micron, the IRS detecting limit for our sample in the rest-frame.  However, the well-known inflection around 1\\micron\\ is also well defined by the red optical spectroscopy and NIR 2MASS photometry.  It is therefore very clear that somewhere between 1 and 4\\micron, there is an infrared bump, which is further supported by the fact that the NIR K-band data point starts to rise toward MIR in both RL and RQ SEDs. Other studies have suggested that there is a 3\\micron\\ bump, resulting from the hottest dust in AGNs \\citep{Netze07,Hiner09}.   Although MSED94 have a lot of upper limits in the FIR from IRAS while we use the latest \\spitzer\\ MIPS data, the SEDs agree surprisingly well for RQ sample in the FIR and extending to the radio.  For the RL sample, it is expected that our \\spitzer\\ data define a better FIR SED, which falls more steeply toward longer wavelengths.  Our radio SED is more luminous than that of MSED94, simply because there are more radio-luminous objects in our sample.  \\subsubsection{Comparison with Quasar SEDs of \\citet{Richa06} }  We compare our SEDs with those of \\citet[][hereafter R06]{Richa06} in Figure~\\ref{fg:compareR06}.  R06 has provided the broadest frequency coverage, from Far-IR to X-ray, in recent SED studies, and their sample of 259 SDSS quasars extends to higher redshift and higher luminosity than ours.  They constructed the SEDs using photometry points, including 5 SDSS magnitudes and 4 \\spitzer\\ InfraRed Array Camera (IRAC) flux densities, supplemented by available \\galex\\ $f$ and $n$ bands, $J$, $H$, and $K$, the \\iso\\ 15\\,\\micron\\ band, and the \\spitzer\\ MIPS 24 and 70\\,\\micron\\ bands.  When objects do not have measurements in the supplemental bands, they used \\citet{Elvis94} SEDs and their ``gap repair'' technique to estimate the missing values. Their X-ray fluxes were obtained from \\rosat\\ detections.  When no detection is available, they estimated the X-ray flux using the 2500\\,\\AA\\ flux and the tight $L_{UV}(2500\\mbox{\\AA})$-$L_{X}$ relationship \\citep{Strat05}.  They have only 8 radio-loud quasars, so the final SEDs are essentially for radio-quiet quasars.  Therefore, we only compare our radio-quiet SED with theirs, but plot both of our RL and RQ SEDs in Figure~\\ref{fg:compareR06} for completeness.  We further choose to compare only with their two SEDs constructed with optically red and blue halves of their sample, and their mean SED would lie between the two.  Our SEDs show relatively much higher X-ray emission, indicating that our sample (and that of \\citet{Elvis94}) is not representative of the SDSS quasars in this region and probably does not follow the $L_{UV}(2500\\mbox{\\AA})$-$L_{X}$ relationship found in SDSS quasars \\citep{Strat05}. In the Far-IR to near-IR region, the overall shape and trend seem to match very well, but ours have more detailed features.  Although the SEDs are normalized at 4200\\,\\AA\\ (Fig.~\\ref{fg:compareR06}), their shapes also match well over most of the UV-optical region.  Only at the two ends, our SED show optical redder and UV brighter.  This implies the different sample properties, because our objects are mostly UV-bright quasars and are probably redder in the $i$ band compared with R06 sample.  Many objects in the R06 sample have much higher redshifts and luminosity than ours, but we do not have enough information to address any possible evolution or luminosity-dependent effects in quasar SEDs by comparing them.  \\subsubsection{Comparison with Other Quasar Composites}  Figure~\\ref{fg:compare} shows the comparison between our SEDs from Far-UV to Mid-IR and other quasar composite spectra, including the HST ultraviolet composites \\citep[][]{Telfe02}, the composite from Sloan Digital Sky Survey \\citep[SDSS,][]{Vande01}, and a Near-IR composite of 27 SDSS quasars \\citep{Glikm06}.  The HST composites extend beyond the FUV part of our SEDs to the extreme-UV (EUV), revealing the EUV peak more clearly.  In the overlapping region, our SEDs are in good agreement with the HST composites although the radio-quiet HST composite drops a little steeper at higher frequency.  The SDSS composite includes both radio-loud and radio quiet objects and is consistent with our SEDs between about 1200 and 4500\\,\\AA.  The increased flux at the red part of the SDSS composite is partly due to host galaxy contamination at low redshift, while the blue part beyond \\lya\\ should be ignored because of \\lya\\ forest contamination.  The NIR composites of radio-loud and radio-quiet samples show little difference in \\citet{Glikm06}.  We use the composite of their entire sample and its optical part matches our SEDs well.  Its NIR region shows the 1\\micron\\ inflection as expected, and the overall shape connects our MIR and optical SEDs surprisingly well.  Although our NIR points are located a little lower than the composite, the continuum trends are consistent.   \\subsection{Uncertainty and Caveats}  As in all previous quasar SED studies \\citep[e.g.,E94; R06][]{}, the dispersion of the mean or composite SEDs is large.  The dispersion of our median SEDs can be evaluated in Fig.~\\ref{fg:msed}, where population standard deviations of all rebinned points in each bin are plotted at the bottom. Because we normalized the individual SEDs at 4215\\,\\AA, the dispersion is minimized in the optical and increases toward both low and high frequencies to about 0.6 (dex) in radio and X-ray.  Even in the NIR to UV region, the standard deviation increases rapidly away from the normalization wavelength.  We also show the dispersion in Figure~\\ref{fg:dispersion}, where we plot all data from normalized SEDs of individual objects and the median SEDs built from them.  The actual difference between the individual SEDs can be more than 2 orders of magnitude.  Even in the NIR to UV, the difference is still more than one order of magnitude.  These all indicate the large variation of quasar SEDs.  We note a caveat that the NIR host galaxy corrections may have a large uncertainty (\\S\\ref{sec:host}, Table~\\ref{tb:hostfrac}), because the H-band host galaxy measurements with ground-based infrared data may be somewhat inherently uncertain.  The host galaxy fractions from HST observations \\citep[][and our own observations]{McLMcL01}are systematically smaller than those from ground-based observations \\citep[][and our IRTF data]{McLRie94a,McLRie94b}.  Two objects have both ground-based and HST H-band observations and show revised lower host galaxy fractions from HST data (Table~\\ref{tb:hostfrac}), with PG~1322+659 from 43\\% to 22\\%, and PG~1427+480 from 45\\% to 26\\%, respectively. In addition, although the same technique was used in estimating the host galaxy fraction with all ground-based data, host galaxy fractions from \\citet[][]{McLRie94a} and our IRTF data seem to be systematically higher than \\citet[][]{McLRie94b} by about 20\\%.  Some SEDs show discrepancy between NIR broad-band points and optical spectra (Fig.~\\ref{fg:allsed}), indicating possible over-subtraction of host galaxy contributions.  Readers should be aware of this issue when using data of this region for their own applications.  Our sample is heterogeneous, and there is the possibility that the sample may be biased.  \\citet{Jeste05} pointed out that BQS quasars are representative of bright blue quasars, but not representative of bright red quasars. Since our sample involves many PG quasars, it is possible that we might be missing some red quasars, and therefore the SEDs are not truly representative of $all$ UV-bright quasars. If this is true, it will affect the spectral index in the optical region. This does not seem to be a serious problem, however, because our SEDs do not show a drop off in the red part of the optical region when compared with the SEDs of R06 (Fig.~\\ref{fg:compareR06}). Moreover, even for the distinguished blue and red quasars in R06, except for the different spectral indices in the optical, their mean SEDs are very similar (see Fig.~\\ref{fg:compareR06}, or their Fig.~11), the largest difference in $log(\\nu f_\\nu)$ is about 0.1 dex, significantly smaller than the dispersion of either R06 or our composite SEDs.  Therefore, the large variation of individual quasar SEDs is still the dominant factor for the uncertainty of the composite SEDs.     }  We have compiled SEDs for 85 quasars using high-quality multi-wavelength data from radio to X-ray energies.  The data were obtained from the next generation space telescopes and ground-based telescopes. Using these data, we have constructed composite (median) SEDs for radio loud and radio quiet quasars.  This work is an update on the mean SEDs built by \\citet{Elvis94} with about twice as many objects. On our website\\footnote{\\url{http://physics.uwyo.edu/agn/}} and the online version of the Journal, we make the electronic version of the median SEDs available for public use.  We caution again that, because of the large variation in quasar SEDs, any composite SEDs should be used with care.  Our composite SEDs are representative only for UV-optical bright quasars.  The RQ median SED is constructed from low-redshift ($z<0.5$) objects, while the RL median SED comes from objects of redshift up to 1.4.  We also plan to investigate the multi-wavelength data of individual objects. We have measured all the spectral parameters of the entire sample and the work will be presented in a separate paper. We will be able to obtain the bolometric luminosities from real data and investigate the bolometric corrections associated with spectral properties, such as continua or emission features. We have also planned to investigate how the quasar SED varies with different physical parameters such as black hole mass and Eddington ratio. The multi-wavelength data will help us to better understand the quasars. A series of papers based on this data set are forthcoming."
    },
    "1107/1107.2121_arXiv.txt": {
        "abstract": "We present our temporal and spectral analyses of 29 bursts from \\sgrnos, detected with the Gamma-ray Burst Monitor onboard the {\\it Fermi} Gamma-ray Space Telescope during the 13 days of the source activation in 2008 (August 22 to September 3). We find that the $T_{90}$ durations of the bursts can be fit with a log-normal distribution with a mean value of $\\sim 123$\\,ms. We also estimate for the first time event durations of Soft Gamma Repeater (SGR) bursts in photon space (i.e., using their deconvolved spectra) and find that these are very similar to the $T_{90}$s estimated in count space (following a log-normal distribution with a mean value of $\\sim 124$\\,ms). We fit the time-integrated spectra for each burst and the time-resolved spectra of the five brightest bursts with several models. We find that a single power law with an exponential cutoff model fits all 29 bursts well, while 18 of the events can also be fit with two black body functions. We expand on the physical interpretation of these two models and we compare their parameters and discuss their evolution. We show that the time-integrated and time-resolved spectra reveal that $E_{\\rm{peak}}$ decreases with energy flux (and fluence) to a minimum of $\\sim30$\\,keV at $F=8.7\\times10^{-6}$ erg cm$^{-2}$ s$^{-1}$, increasing steadily afterwards. Two more sources exhibit a similar trend: SGRs J$1550-5418$ and $1806-20$. The isotropic luminosity, $L_{\\rm iso}$, corresponding to these flux values is roughly similar for all sources ($0.4-1.5\\times10^{40}$ erg s$^{-1}$). ",
        "introduction": "}  Magnetars are slowly rotating neutron stars associated with extreme magnetic fields ($B>10^{14}$ G). Several obscure neutron star subpopulations haven been claimed as magnetar candidates, in particular Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs); for reviews see \\citet{woods2006,mereghetti2008}. Most magnetars have been discovered either from their persistent X-ray emission properties or when they enter into randomly occurring outbursts, during which they emit a multitude of short ($\\sim100$ milliseconds), soft $\\gamma-$/hard X$-$ray bursts. Thus far, approximately twenty magnetar sources are known and most of them reside in the Galactic Plane with a higher concentration close to the center; two are located in the Magellanic Clouds.  \\sgr was discovered with \\textit{Swift} on 2008 August 22, when it emitted a series of bright, soft, short bursts \\citep{holland2008,barthelmy2008}. The first burst also triggered the Gamma-ray Burst Monitor (GBM) onboard the {\\it Fermi} Gamma-ray Space Telescope ({\\it FGST}; hereafter {\\it Fermi}). Soon after, our Target of Opportunity (ToO) observations with the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) established a period of $\\sim$ 5.76 s in the persistent X-ray emission of the source \\citep{gogus2008}. Further observations with \\textit{RXTE} and the \\textit{Swift}/X-ray Telescope (XRT) revealed a spin-down rate of $1.5(5)\\times10^{-11}$ s s$^{-1}$, indicating a dipole magnetic field of $2.0\\times10^{14}$ G \\citep{woods2008, rea2009, gogus2010}. Our subsequent $\\textit{Chandra}$ ToO observations established an accurate location of the source at $R.A. (J2000) = 05^{\\rm h} 01^{\\rm m} 06\\fs76$, $Dec (J2000) = +45\\arcdeg 16\\arcmin 33\\farcs92$, with an 1$\\sigma$ uncertainty of $0\\farcs11$ \\citep{gogus2010}. This is the first magnetar location at roughly the Galactic anticenter direction, placing \\sgr most likely at the Perseus arm of our Galaxy at $\\sim2$ kpc \\citep{xu2006}.  The \\sgr outburst lasted approximately 2 weeks, during which several bright bursts were detected with \\textit{Swift}, GBM, {\\it RXTE}, {\\it KONUS}-Wind, and {\\it Suzaku} (Enoto et al. 2009; Aptekar et al. 2009; Kumar et al. 2010; see also Table \\ref{obs}). After the burst activity ceased, the source flux decreased exponentially with an e-folding time of $27.9\\pm2.5$ days \\citep{gogus2010}. During the entire outburst interval, GBM triggered on 26 bursts; in addition, an untriggered event search in the daily data sets revealed seven more bursts. We present here our analyses of the 29 GBM bursts from \\sgr for which we have high-resolution data; the properties of the X-ray persistent source emission have been published by \\citet{rea2009} and \\citet{gogus2010}. In Section \\ref{data}, we describe the instrument, observations and data selection. In Sections \\ref{tempo} and \\ref{spec}, we present our temporal and spectral results, respectively. We discuss the interpretation of our results in Section \\ref{discussion}. ",
        "conclusions": "} \\subsection{COMPT {\\it versus} BB+BB} Our spectroscopic fitting clearly indicates that the COMPT and BB+BB models yield superior fits to the other possibilities. The COMPT model, with its power-law shape curtailed by an exponential turnover, is intended to mock-up the classic unsaturated Comptonization spectrum realized in models of accretion disks such as in Cyg X-1 or in active galactic nuclei (see Chapter 7 of Rybicki \\& Lightman 1979 for a summary of its development as a solution of the Kompaneets equation).  These models use hot, thermal electrons in a corona to repeatedly scatter low-energy photons, heating them gradually up to an energy $E\\sim kT_e$ consistent with the electron temperature $T_e$, at which point a quasi-exponential spectral turnover emerges as further heating becomes impossible. The power law marks the scale-independence of the Compton upscattering, and its slope depends only on the mean energy gain per collision ($\\langle \\Delta E\\rangle =4 kT_e$ for non-relativistic electrons) and the probability of loss of photons from the scattering zone, being expressed as \\begin{equation} \\frac{dN}{dE}\\;\\propto\\; E^{\\lambda} \\quad ,\\quad -\\lambda\\; =\\; -\\frac{1}{2} + \\sqrt{\\frac{9}{4} + \\frac{4}{y}}\\quad , \\label{Comp_index} \\end{equation} with the Compton $y$-parameter in the domain $y< 1$. This parameter is the product of the average fractional energy change per scattering and the mean number of scatterings, and is $4 kT_e/(m_ec^2)\\, \\hbox{max} \\{\\tau,\\, \\tau^2\\}$ for a non-relativistic situation for repeated Compton upscattering.  Here $\\tau$ is the scattering optical depth.  In the context of magnetars, such Comptonization can also ensue, but the strong magnetic field now plays an important role.   Ephemeral coronae of hot electrons could be expected in a dynamic inner magnetosphere. For example, this could be due to intense dissipation of magnetic energy in the closed field line region via field line twisting, i.e., transient departures from poloidal geometry, as in the considerations of Thomson, Lyutikov \\& Kulkarni (2002), Thompson \\& Beloborodov (2005),  Beloborodov \\& Thompson (2007), and Nobili, Turolla \\& Zane (2008).  Such a pumping of energy into electrons in low altitude regions would then be subject to irradiation by the intense bath of surface X-ray emission. The electron coronae would mimic those of their black hole counterparts discussed above, and serve as a Comptonization target for the X-rays. Temporally, the coronae could be quite variable, resulting in varying or chaotic time profiles such as are observed.  If the deposition of energy in hot electrons is persistent over many light-crossing timescales, so also can the upscattered hard X-ray emission be.  In this magnetic case, the spectrum would again naturally be a power law truncated by an exponential tail, whose energy is pinned by $T_e$.  The slope does differ somewhat in value from the non-magnetic case because the presence of the intense field anisotropizes the environment.  This alters the average fractional energy change per scattering from the isotropic $4 kT_e$ form in Rybicki \\& Lightman (1979). We note that this picture is similar to that of Lyutikov \\& Gavriil (2006; see also Rea et al. 2008), who modeled the steep X-ray tails below 10 keV in quiescent magnetar emission using a resonant cyclotron scattering picture.  Since the electrons will move along the field lines in the zeroth Landau level, the exact kinematics that impact the determination of $\\langle\\Delta E\\rangle$ depend on the colatitude and altitude of the collisions.  Therefore, the effective {\\it magnetic Compton $y$-parameter}, $y_B$, that could be substituted in Eq.~(\\ref{Comp_index}), would take a somewhat different value from its non-magnetic cousin, but one anticipates that the range of spectral realizations would be similar to the $B=0$ case.  Since head-on collisions generally yield greater heating of photons, and these are more likely at lower altitudes, it is expected that interaction zones close to the stellar surface would spawn larger $y_B$ and therefore flatter photon spectra.  If there were a coronal radius expansion, this would then translate to an evolutionary steepening of the spectra with time.  Note that the photon retention probabilities in the coronae will need to be higher than in the Lyutikov \\& Gavriil (2006) scenario to generate the requisite moderate optical depths $\\tau$ to match the flat outburst spectra discussed in this paper.  The scatterings will take place below the cyclotron resonance unless their altitude is above around 10 stellar radii (e.g. Lyutikov \\& Gavriil 2006).  Below resonance, the cross section is far inferior to the Thomson value (e.g. see Herold 1979).  Then, the coronal electron density must accordingly be much higher than when $B=0$ in order to establish a sizable optical depth. This provides a possible distinction between the two classes of magnetar emission: the steady $<10$\\,keV signal may originate at higher altitudes where the resonance is accessed and the local electron density is low, whereas the bursts reported here may be triggered closer to the surface in higher density zones that initially precipitate scattering below the cyclotron resonance. As photons get energized above 20 keV, the altitude where the resonance is accessed is lowered, increasing the cross section for scattering and therefore $y_B$.  Note also, that to a considerable extent, polarization mode-switching (e.g. Miller 1995) between the two photon polarizations can help increase the opacity, adding a further nuance to consider when modeling the scattering environment.  Observationally, it is difficult to discern unambiguously the presence of the intense field using a Comptonization model scenario in the energy range of data presented in this paper: details of both the field strength and the geometry are subsumed in a single parameter $y_B$. Since the emission in this scenario should be strongly polarized, and the degree of polarization should depend on the interaction geometry, a hard X-ray polarimeter would provide insightful probes into the presence of a strongly anisotropizing super-critical field.  The dual blackbody can also be envisaged as a viable alternative from a theoretical standpoint.  The moderate Thomson depths required to generate flat Comptonization spectra could easily be higher, thereby pushing the electron-photon interactions more towards an equilibration. The saturation temperature is then controlled by the largely uncertain total energy dissipated in the inner magnetosphere per hot electron present. Since any equilibration will be non-uniform over a coronal volume, there should be a modest temperature gradient throughout, smearing out the continuum.  The range of temperatures will not be great because the system is not gravitationally hydrostatic in character.  Moreover, radiative transfer effects impact the spectral shape and further modify it.  Accordingly, pure blackbody shapes are not expected.  It is quite conceivable that a two-component blackbody fit may well approximate the emergent continuum that is a superposition of distorted blackbodies spanning a small range of temperatures.  Most probably, due to general thermodynamic considerations, the base of the coronae (e.g. centered near the source of magnetic dissipation) should be hotter than the outer layers (see also Lyubarsky 2002 and Thompson \\& Duncan 1995).  Interestingly, the fits here generate a smaller volume associated with the hotter blackbody contribution, consistent with the expectation for coronal structure.  Yet, somehow, we must obtain a view of the hotter zone, therefore indicating a strongly aspherical coronal geometry.  Modeling this semi-equilibration is a challenging task for theorists considering the influences of the field, the twisted magnetospheric geometry, and the inherent anisotropy and polarization-dependence of the scattering process. At present, it is not possible to distinguish between this thermal scenario, and a Comptonization one.  \\subsection{Conclusions} Since the COMPT model fits all events, we used it to derive fluences (time-integrated and time-resolved), and determined -- using the time-resolved values -- that the hardness-fluence correlation can be described by a broken power law, with a minimum at a fluence of $1.7\\times10^{-7}$ erg cm$^{-2}$. We have used $E_{\\rm peak}$ values to characterize the hardness of the events in our sample. Earlier studies of this correlation used hardness ratios to bypass proper spectral fits due to low quality or insufficient data. Moreover, these studies \\citep{fenimore1994,gogus2001} had a very narrow overlap in fluence space. \\citet{fenimore1994} used 95  SGR\\,$1806-20$ events detected with the Interplanetary Sun Earth Explorer-3 (ISEE-3) with fluences ranging between $1.25\\times10^{-7} - 8\\times10^{-6}$ erg cm$^{-2}$, and found a slightly positive correlation between the two quantities (hardness increasing with fluence). On the other hand, \\citet{gogus2001} analyzed 159 and 385 events from SGRs $1806-20$ and $1900+14$, respectively, observed with {\\it RXTE}, with fluences ranging between $1.0\\times10^{-9} - 2.0\\times10^{-7}$ erg cm$^{-2}$, and found the opposite trend (an anticorrelation between hardness ratios and fluences). Since the GBM data sample covers both fluence ranges ($2.0\\times10^{-8} - 2.0\\times10^{-5}$ erg cm$^{-2}$) we were able to establish that both trends are indeed correct, and to define the turning point in the $E_{\\rm peak}$ - fluence diagram.  Only two magnetar candidates were observed with GBM to emit a multitude of bursts, thus allowing us to construct their $E_{\\rm peak}$ - flux diagrams. We find for the time-resolved data of \\sgrnos, that $E_{\\rm peak}$ reaches a minimum of $\\sim30$\\,keV at a flux value of $8.7\\times10^{-6}$ erg cm$^{-2}$ s$^{-1}$. The second source (SGR\\,J$1550-5418$; van der Horst et al. 2011 in preparation) has flux values of $4.4\\times10^{-6}$ erg cm$^{-2}$ s$^{-1}$ at a similar $E_{\\rm peak}$ minimum. In addition we used the hardness ratio - count rate relationship of SGR\\,$1806-20$ found in the {\\it Integral} data \\citep{gotz2004} to derive an approximate flux value of $1.7\\times10^{-6}$ erg cm$^{-2}$ s$^{-1}$ (using their conversion from count rate to flux) at the hardness ratio minimum \\citep[Figure 3 in][]{gotz2004}. We then converted these values into isotropic source luminosity, $L_{\\rm iso}$ (see Table \\ref{liso}); we note that (with the caveat of small number statistics, and uncertainties in the distance measurements) these values are comparable [$(0.4-1.5)\\times10^{40}$ erg s$^{-1}$], although the $B-$fields and fluxes at $E_{\\rm peak}$ minima vary by approximately a factor of ten between the two GBM sources and SGR\\,$1806-20$. Whether these differences reflect intrinsic source properties or instrumental effects is not yet clear. The $E_{\\rm peak}$ trend is, however, clearly established in at least two different instrument data sets. The physical interpretation of this trend is beyond the scope of this paper.  Part of our sample (18 bursts) was fit with the BB+BB spectral function applied to time-integrated intervals; the remaining 11 events could not be fit due to poor statistics. In addition for five bright events we performed time-resolved spectral analysis. We found that the temperatures, emission areas and fluences (luminosities) of the two thermal components exhibit very similar properties in both cases (time-integrated and -resolved). Our results are consistent with those presented for intermediate and short bursts from SGR $1900+14$ by \\citet{olive2004} and \\citet{israel2008}; the \\sgr burst emission areas and temperatures fall into the same region with those of the short bursts of SGR $1900+14$ shown in Figure 5 of \\citet{israel2008}. We also see a similar trend in the correlation between the luminosities of the two components as described by \\citet{israel2008}, namely that both luminosities increase in tandem. This behavior indicates that the hot and cool BB components may come from two separate emission regions, as also pointed out by \\citet{israel2008}: a smaller but hotter one from the surface of the magnetar, and a larger, cooler one from the star's magnetosphere (but see also the discussion in section 5.1). We note here that the time-integrated hot BB emission area of 13 \\sgr bursts (i.e., not including the 5 brightest events; Figure \\ref{r2_kt}) is similar to the emission area ($\\sim 0.05$ km$^2$) of the BB component found in the persistent emission of SGR J1550$-$5418 during one active bursting episode in January 2009, believed to originate from a hot spot on the neutron star surface \\citep{kaneko2010}.  Finally, we performed a detailed temporal analysis of all 29 bursts and estimated their durations ($T_{90}$/$T_{50}$) for the first time in count {\\it and} in photon space. Both estimates agree within statistics and thus validated all earlier (count space) duration estimates of magnetar candidate bursts. The durations (and for four SGRs also the emission times and duty cycles) of five magnetar candidates follow a log-normal or normal distribution. We find that \\sgr events are very similar in average duration with four more magnetar candidates (three SGRs and one AXP). However, the $T_{90}$ distribution of the bursts from AXP\\,1E$2259+586$ has a factor of two larger dispersion \\citep[$\\sigma \\sim 0.73$,][]{gavriil2004} than those of at least three other SGRs \\citep[$\\sigma \\sim0.34, 0.35$,][ and $\\sim0.35$, present work]{gogus2001}. \\sgr bursts have an average duty cycle ($\\delta_{90}=0.68$) larger than other SGRs \\citep[0.45, 0.46 for SGRs $1900+14$, $1806-20$, respectively;][]{gogus2001}. The differences in the intrinsic properties of the sources may be due to differences in the size of the active region responsible for the burst emission. We discuss below the burst energetics and its evolution. Similarly to other magnetars, we also do not find a correlation between the pulse phase of the source and the burst peak times.  The overall behavior of \\sgr during its 13-day active period (2008 August 22 to September 3) is also very interesting.  Although about half of the bursts (16/29) occurred on one day (2008 August 23), most of the burst energy was emitted from the source during August $24-26$ (see also Figure \\ref{energetic}). The average fluence of the beginning and the later part of the active period was constant and at a lower level. During these three days, GBM detected seven bursts, five of which are also distinguished in that the emission areas from both BB components are the largest. Further, one of these five events (bn080824.054) shows a double-peaked structure reminiscent of some bright thermonuclear Type I X-ray bursts from accreting neutron stars, which has been interpreted as due to photospheric radius expansion (PRE). \\citet{watts2010} studied the effects of PRE in high magnetic fields using this event and find that the predicted flux from PRE theory is consistent with the one observed, opening the way to determining fundamental parameters of neutron stars, such as their equation of state."
    },
    "1107/1107.2067_arXiv.txt": {
        "abstract": "The detection of narrow SiO thermal emission toward young outflows has been proposed to be a signature of the magnetic precursor of C-shocks. Recent modeling of the SiO emission across C-shocks predicts variations in the SiO line intensity and line shape at the precursor and intermediate-velocity regimes in only few years. We present high-angular resolution (3.8$''$ $\\times$3.3$''$) images of the thermal SiO $J$=2$\\rightarrow$1 emission toward the L1448-mm outflow in two epochs (November 2004-February 2005, March-April 2009). Several SiO condensations have appeared at intermediate velocities (20-50$\\,$km$\\,$s$^{-1}$) toward the red-shifted lobe of the outflow since 2005. Toward one of the condensations (clump D), systematic differences of the dirty beams between 2005 and 2009 could be responsible for the SiO variability. At higher velocities (50-80$\\,$km$\\,$s$^{-1}$), SiO could also have experienced changes in its intensity. We propose that the SiO variability toward L1448-mm is due to a real SiO enhancement by young C-shocks at the internal working surface between the jet and the ambient gas. For the precursor regime (5.2-9.2$\\,$km$\\,$s$^{-1}$), several narrow and faint SiO components are detected. Narrow SiO tends to be compact, transient and shows elongated (bow-shock) morphologies perpendicular to the jet. We speculate that these features are associated with the precursor of C-shocks appearing at the interface of the new SiO components seen at intermediate velocities. ",
        "introduction": "Variability is a common phenomena in the evolution of very young stars. \\citet{joy45} very early found that T Tauri stars have an irregular behavior in their optical emission. More recently, simultaneous X-ray and optical/near-IR variability studies have shown that a high fraction of pre-main-sequence stars are variable in X-ray and optical/near-IR emission \\citep{sta06,for07}. At radio wavelengths, young stellar objects are known to show variability in their radio continuum emission in time-scales of few years \\citep[][]{cur06,pech10}. In the case of molecular line emission, only {\\it non-thermal} maser emission has been found to be variable \\citep[see][]{reid81,eli92}. However, variability in the morphology, line intensity and line profiles of {\\it thermal} molecular emission still remains to be reported.  Broad SiO thermal line emission is known to be a good indicator of shocked gas in molecular outflows \\citep[][]{mar92}. The SiO line profiles have typical linewidths of tens of km$\\,$s$^{-1}$ and appear shifted with respect to the ambient cloud velocity. In addition to this broad emission, \\citet{lef98} and \\citet{jim04} reported the detection of very narrow SiO thermal lines (linewidths of $\\leq$1$\\,$km$\\,$s$^{-1}$) centered at almost ambient velocities toward the young NGC1333 and L1448-mm outflows. Toward L1448-mm, \\citet{jim04} proposed that narrow SiO could arise from the magnetic precursor region within very young C-shocks \\citep{jim04}.  At the early stages of the C-shock propagation, the magnetic precursor of C-shocks accelerates, compresses and heats the ion fluid before the neutral one \\citep[][]{dra80}. The subsequent ion-neutral velocity drift generates enough friction to sputter dust grains \\citep{cas97,gus08}, progressively enhancing the gas-phase abundances of SiO \\citep{jim08}. \\citet{jim09} have claimed that the SiO line profiles in young outflows should then evolve from very narrow lines arising from the magnetic precursor region, to broader SiO emission generated in the intermediate- and high-velocity postshock gas. From this, variability of the narrow SiO lines and of the intermediate-velocity SiO emission should appear in only few years \\citep{jim09}.  We present interferometric images of the SiO $J$=2$\\rightarrow$1 thermal line emission toward the L1448-mm outflow observed with the PdBI at two different epochs (2005 and 2009)\\footnote{Based on observations carried out with the IRAM Plateau de Bure Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain)}. These images show that several SiO condensations (named A, B, C, D and E) have recently appeared at intermediate velocities along the direction of the red-shifted lobe of the outflow since 2005. The increase in the SiO flux toward condensations A, B and C are likely real and produced by an enhancement in the gas-phase abundance of SiO by the recent interaction of young C-shocks with the ambient gas. The variable nature of the SiO clumps D and E, however, remains unclear (Section$\\,$\\ref{mod}). Our results also reveal the detection in 2009 of several new narrow SiO components that tend to be found at the interface between the SiO condensations A, B, C, D and E reported at intermediate velocities toward the red-shifted lobe of the outflow (Section$\\,$\\ref{pre}). We speculate that the recently detected narrow SiO components are possibly associated with the magnetic precursor of C-shocks (Section$\\,$\\ref{origin}). This is the first time that variability of molecular thermal emission is ever reported. ",
        "conclusions": "\\label{con}  In this paper, we present interferometric images of the thermal emission of the SiO $J$=2$\\rightarrow$1 rotational line observed toward the L1448-mm outflow at two different epochs (2005 and 2009). The PdBI images reveal that this emission is variable in only 4$\\,$yrs. Not only the morphology of the SiO emission, but also its line profiles, have significantly changed since 2005, with a clear increase of the SiO flux at intermediate-velocities, and possibly at high-velocities. This suggests that new material has recently entered the shock.  For the precursor component, we have detected several narrow SiO velocity components (M1, M2, M3 and M4) with an elongated morphology in the direction perpendicular to the high-velocity jet. These features resemble those found in bow-shocks. Although these components could have appeared recently, new multi-epoch and higher sensitivity interferometric images are needed to clearly establish the variable nature of the narrow SiO emission detected toward L1448-mm.  We speculate that the variability of the thermal SiO $J$=2$\\rightarrow$1 emission observed toward the L1448-mm outflow could be explained by the time-dependent propagation of very young C-shocks (with their precursors) through the ambient molecular cloud. Monitoring of the variability of the thermal SiO emission toward young molecular outflows could be a useful tool to study the evolution of these objects, and their impact on the surrounding medium."
    },
    "1107/1107.2298_arXiv.txt": {
        "abstract": "{The large Coalsack dark cloud is one of the most prominent southern starless clouds, which is even visible to the naked eye. Furthermore, it is one of the rare molecular clouds without clear signs of star formation.} {We investigate the dynamical properties of the gas within the Coalsack.} {The two highest extinction regions were mapped with the APEX telescope in $^{13}$CO(2--1) comprising a region of $\\sim$1\\,square degree.} {In addition to the well-known, nearby gas component around $-4$\\,km\\,s$^{-1}$, we identified additional molecular gas components -- in particular a second extended molecular cloud at a velocity of $\\sim -30$\\,km\\,s$^{-1}$ and an estimated distance of $\\sim$3.1\\,kpc -- that dominate the column density and visual extinction distributions in the northeastern part of the Coalsack. Although comprising $\\sim$2600\\,M$_{\\odot}$, the mass of this distant cloud is distributed over an extent of $\\sim$73\\,pc, much larger than typical high-mass infrared dark clouds. Its filamentary structure is consistent with a compressible gaseous self-gravitating cylinder, and its low mass per length indicates that it may be stable against gravitational collapse.  We find barely any mid-infrared emission in archival MSX data, which is indicative of almost no star-formation activity in the near and far cloud complexes. The nearby clouds have narrow, almost thermal velocity dispersions with median values between 0.2 and 0.4\\,km\\,s$^{-1}$, which is also consistent with low star-formation activity. Only Tapia's Globule 2 exhibits a velocity dispersion increase toward the extinction peak and peak-velocity gradients over the core, which is indicative of a state of elevated dynamical properties.} {The Coalsack is not one single coherent structure, but consists of several cloud complexes nearby as well as at several kpc distance. All studied clouds appear as starless low-turbulence regions that may not even collapse in the future. Only one globule exhibits more dynamical signatures and is a good candidate for present/future star formation.} ",
        "introduction": "\\label{intro}   The Coalsack is one of the closest and most widely-recognized dark clouds in the southern hemisphere (between Galactic longitudes $305\\geq l \\geq 300$\\,deg) because it is visible to the bare eye (e.g., \\citealt{nyman2008}).  \\citet{kainulainen2009} recently produced extinction maps based on 2MASS near-infrared data of all supposedly nearby molecular clouds including the Coalsack. Figure \\ref{extinction} shows the large-scale extinction ($A_v$) map of this region. While the visual extinction through the Coalsack is generally low, there are a few regions where the extinction $A_v$ rises above 6 and even 10\\,mag. Although the extinction maps set significant constraints on the morphology and column density distributions, allowing us to derive results about cloud structure and stability as well as potential column density star-formation thresholds (\\citealt{kainulainen2009,kainulainen2011}), these data lack any kinematical information about the cloud. Furthermore, this cloud is one of the rare cases of a nearby molecular cloud devoid of any star-formation signatures (e.g., \\citealt{nyman2008,kainulainen2009}). This is one of the reasons why it has not received as much attention in the past as typical star-forming clouds like Taurus of $\\rho$ Ophiuchus. Despite this, the Coalsack therefore constitutes an excellent laboratory to study pristine and undisturbed molecular clouds.  \\begin{figure}[htb] \\includegraphics[angle=-90,width=90mm]{extinction.eps} \\caption{2MASS extinction map of the whole Coalsack by \\citet{kainulainen2009}. The contour levels are at 2, 6 and 10\\,mag extinction. The two boxes outline the two regions mapped in $^{13}$CO(2--1) and C$^{18}$O(2--1) for this work.} \\label{extinction} \\end{figure}  Earlier spectral-line mapping on the Coalsack was conducted in $^{12}$CO by \\citet{nyman1989} at a spatial resolution of 8.8$'$, and in $^{13}$CO by \\citet{kato1999} at a resolution of 2.7$'$. Thus, even basic dynamical measures (e.g., central velocities, dispersions, etc.) were limited to spatial scales larger than what is expected for molecular clumps or cores (or equally limited in dynamical range/velocity resolution).  The typical velocities of the region are between -6 and 0\\,km\\,s$^{-1}$ (e.g., \\citealt{nyman1989}), and the distance estimates vary between 100 and 200\\,pc (e.g., \\citealt{franco1995,corradi1997,knude1998,nyman2008}). \\citet{nyman1989} reported the presence of an additional velocity component at $-35$\\,km\\,s$^{-1}$, which they attribute to a background cloud associated with the Sagittarius-Carina arm that is unlikely to be responsible for any of the extinction features.  Some individual fields of the Coalsack have been mapped in tracers of higher density by some authors, mainly concentrating on the region known as ``Tapia's Globule 2'', which presumably is one of the few dense cores in the region (e.g., \\citealt{tapia1973,lada2004,rathborne2009b}). Furthermore, \\citet{fontani2005}, \\citet{beltran2006} and \\citet{miettinen2010} studied a smaller sub-filament just adjacent to the eastern edge of our northeastern field. These studies, however, cover only a minuscule portion of the complex and therefore are not very informative for the dynamics of the complex itself.  \\begin{figure}[htb] \\includegraphics[width=90mm]{spec_aver_13co.eps} \\caption{Average $^{13}$CO(2--1) spectra toward the northeastern region R1 (gray with this histogram) and the southwestern region R2 (thick histogram multiplied by 0.2 to be on the same scale).} \\label{aver} \\end{figure}  To fill this missing link was one of the motivations mapping parts of the Coalsack in spectral line emission. For this purpose, we selected the two major regions with significant fractions of cloud structure above $\\sim$6\\,mag extinction, marked by the two boxes in Figure \\ref{extinction}. The southwestern region also comprises ``Tapia's Globule 2''. These two regions cover in total about 1 square degree and were mapped with APEX in $^{13}$CO(2--1) and C$^{18}$O(2--1) at a spatial resolution of $\\sim 27.5''$. ",
        "conclusions": "Mapping approximately 1 square degree around the regions of the highest extinction in the almost entirely starless Coalsack dark cloud in $^{13}$CO(2--1) with APEX revealed several interesting properties:  \\begin{itemize}  \\item The Coalsack is not one coherent dark low-mass cloud at distances below 200\\,pc and typical velocities of $\\sim -4$ to 0\\,km\\,s$^{-1}$, but we find a second velocity component at $\\sim -30$\\,km\\,s$^{-1}$ corresponding to a second cloud at a distance of $\\sim 3.1$\\,kpc. This second more distant cloud dominates the dust-extinction signatures in the northeast of the Coalsack.  \\item Mid-infrared MSX data reveal barely any star-formation activity toward any of the mapped clouds.  \\item Although the total observed mass of the distant cloud with $\\sim$ 2600\\,M$_{\\odot}$ is far larger than those of the two mapped nearby clouds (on the order of several solar masses), we cannot consider it as a potential high-mass star-forming region, because its mass is distributed over a spatial extent in east-west direction of $\\sim$73\\,pc. Hence, it is smoothly distributed over a large spatial area.  \\item The more distant cloud resembles in shape a large and twisted filament. Analyzing this structure within the framework of an isothermal gravitationally bound cylinder, we find that the data are consistent with a compressible gaseous filament and inconsistent with an incompressible fluid. Furthermore, the characteristic mass per length $M/l$ is low with $\\sim$36\\,M$_{\\odot}$pc$^{-1}$. Hence, this structure is potentially stable against gravitational collapse.  \\item The nearby low-mass clouds both exhibit narrow velocity distributions with median values between 0.2 and 0.4\\,km\\,s$^{-1}$, also indicative of almost no star-formation activity.  \\item Only Tapia's Globule 2 appears different. While the peak velocity dispersion barely exceeds 0.6\\,km\\,s$^{-1}$, we nevertheless see a line-width increase from the map edge toward the extinction peak. Furthermore, an analysis of the central-velocity-increments reveals significant velocity changes at the Globule's edges. Both features are suggestive of dynamical action in that region, e.g., maybe early infall activity.  \\end{itemize}  In summary, our data reveal different cloud components nearby as well as far away. While all observed clouds are mostly bare of any star-formation activity, we find filamentary structures consistent with compressible self-gravitating gaseous cylinders as well as one Globule in an apparent elevated dynamical state."
    },
    "1107/1107.2303_arXiv.txt": {
        "abstract": "{} {We present a number of improvements to the MILES library and stellar population models. We correct some small errors in the radial velocities of the stars, measure the spectral resolution of the library and models more accurately, and give a better absolute flux calibration of the models.} {We use cross-correlation techniques to correct the radial velocities of the offset stars and the penalised pixel-fitting method, together with different sets of stellar templates, to re-assess the spectral resolution of the MILES stellar library and models. We have also re-calibrated the zero-point flux level of the models using a new calibration scheme.} {The end result is an even more homogeneously calibrated stellar library than the originally released one, with a measured spectral resolution of $\\sim$2.5~\\AA, almost constant with wavelength, for both the MILES stellar library and models. Furthermore, the new absolute flux calibration for the spectra excellently agrees with predictions based on independent photometric libraries.} {This improved version of the MILES library and models (version 9.1) is available at the project's website (http://miles.iac.es).} ",
        "introduction": "The MILES library consists of 985 stars spanning a large range in atmospheric parameters (S\\'anchez-Bl\\'azquez et al. 2006, hereafter S06). The spectra were obtained at the 2.5m Isaac Newton Telescope at the Observatorio del Roque de los Muchachos and cover the range 3525-7500\\AA\\ with an originally estimated spectral resolution of 2.3~\\AA\\ full-width half maximum (FWHM). A homogenized compilation of the stellar atmospheric parameters (T$_{\\rm eff}$, Log(g), [Fe/H]) for the stars of this library was presented in Cenarro et al. (2007). Four years after its release, MILES has proved to be an often used, well calibrated stellar library that has been extensively used for a wide range of different applications (Davis et al. 2007; Emsellem et al. 2007; Shields et al. 2007; Martins \\& Coelho 2007; Cenarro et al. 2008; Scelsi et al. 2008; Zuckerman et al. 2008; Cid Fernandes \\& Gonz{\\'a}lez Delgado 2010). The stellar library is also the basis for our empirical models for stellar population synthesis (Vazdekis et al. 2010, hereafter V10). Both the stellar library and models can be easily accessed and manipulated through our dedicated website: http://miles.iac.es.  During the last year, we have had some private communications with different research groups reporting that some stars in the stellar library showed radial velocities different from zero. We have explored this question and, indeed, confirmed these offsets. In this research note, we summarize our solution for this problem and also the re-assessment of the spectral resolution for the stellar library and models, which had also been questioned. While sorting out these concerns, we also addressed the absolute (but not the relative) flux scaling for our model spectra. Updated versions of both datasets are available from our website.  \\begin{figure*} \\begin{center} \\includegraphics[angle=0,width=\\linewidth]{velocity_offsets.eps} \\caption{Histogram with the mean velocity difference (averaged over the whole wavelength range) for stars between the originally released MILES stellar library (v9.0) and the corrected version presented here (v9.1). The dispersion of the histogram ($\\sigma_{\\rm \\Delta V}$) is also indicated. Offsets are also presented as a function of atmospheric parameters: effective temperature (Log(T$_{\\rm eff}$)), surface gravity (Log(g)), and metallicity ([Fe/H]).} \\label{fig:histogram} \\end{center} \\end{figure*} ",
        "conclusions": "In this report we summarize our efforts to improve on a number of minor problems identified in both the MILES stellar library and models. We have corrected some stars in the stellar library for radial velocity offsets with respect to the solar rest frame and checked that they have a small impact on the resulting line-strength indices of the stars and models. We also re-assessed the spectral resolution of the MILES library and models, confirming that it is quite constant as a function of wavelength. We established that the new values are $2.50\\pm0.07$~\\AA\\ for the stellar library and $2.51\\pm0.07$ for the models. We also changed the scheme used to flux calibrate the spectra in the absolute sense. As a result we find an excellent agreement ($\\leq0.01$ mag) between fluxes measured from our spectra and those from predictions obtained by using photometric (rather than spectroscopic) libraries. Our tests also show that the spectral resolution of the Indo-US library is slightly higher (FWHM$=1.36\\pm0.06$~\\AA) than the approximate value established by the Indo-US team (FWHM$\\approx1.2$~\\AA).  The updated version of the MILES stellar library and models in this research note (version 9.1) supersede those presented in S06 and V10 and are available at the MILES website (http://miles.iac.es). This new set of spectra sets the basis for forthcoming papers describing the extension of our model predictions to other wavelength ranges.\\looseness-2 \\begin{figure*} \\begin{center} \\includegraphics[angle=90,width=\\linewidth]{miles_ssp_ls_tests.eps} \\caption{Line-strength differences between versions 9.0 and 9.1 of the MILES SSP models with a Kroupa Universal IMF. Solid line marks the zero level. Dashed and dotted lines indicate the expected uncertainty for spectra with a signal-to-noise ratio of 30 and 75~\\AA$^{-1}$, respectively (Cardiel et al. 1998). The vertical axis was adjusted to twice the maximum uncertainty expected for a spectrum with a signal-to-noise ratio of 30~\\AA$^{-1}$ within the line-strength index range shown.} \\label{fig:miles_ls_ssp} \\end{center} \\end{figure*}"
    },
    "1107/1107.2904_arXiv.txt": {
        "abstract": "In the effective-one-body (EOB) approach the dynamics of two compact objects of masses $m_1$ and $m_2$ and spins $\\mathbf{S}_1$ and $\\mathbf{S}_2$ is mapped into the dynamics of one test particle of mass $\\mu = m_1\\,m_2/(m_1+m_2)$ and spin $\\mathbf{S}_*$ moving in a deformed Kerr metric with mass $M =m_1+m_2$ and spin $\\mathbf{S}_\\mathrm{Kerr}$. In a previous paper we computed an EOB Hamiltonian for spinning black-hole binaries that (i) when expanded in post-Newtonian orders, reproduces the leading order spin-spin coupling and the leading and next-to-leading order spin-orbit couplings for any mass ratio, and (iii) reproduces {\\it all} spin-orbit couplings in the test-particle limit. Here we extend this EOB Hamiltonian to include next-to-next-to-leading spin-orbit couplings for any mass ratio. We discuss two classes of EOB Hamiltonians that differ by the way the spin variables are mapped between the effective and real descriptions. We also investigate the main features of the dynamics when the motion is equatorial, such as the existence of the innermost stable circular orbit and of a peak in the orbital frequency during the plunge subsequent to the inspiral. ",
        "introduction": "\\label{sec:intro}  Coalescing compact binaries composed of neutron stars and/or black holes are among the most promising gravitational-wave sources for ground-based detectors, such as the Laser Interferometer Gravitational-wave Observatory (LIGO)~\\cite{Abbott:2007}, Virgo~\\cite{Acernese:2008}, GEO~\\cite{Grote:2008zz}, the Large Cryogenic Gravitational Telescope (LCGT)~\\cite{Kuroda:2010}, and future space-based detectors.  So far, the search for gravitational waves with LIGO, GEO and Virgo detectors has focused on non-spinning compact binaries~\\cite{Abbott:2005kq,Abbott:2007xi, Abbott:2009qj,Abadie:2010yba,Abadie:2011kd}, although in Ref.~\\cite{Abbott:2007ai} single-spin templates were used to search for inspiraling spinning compact objects. Within the next 4-5 years LIGO and Virgo detectors will be upgraded to a sensitivity such that event rates for coalescing binary systems will increase by a factor of one thousand. Thus, it is timely and necessary to develop more accurate templates that include spin effects. For maximally spinning objects, we expect that reasonably accurate templates would need to be computed at least through 3.5PN order. In the non-spinning case, studies at the interface between numerical and analytical relativity have demonstrated that templates computed at 3.5PN order are indeed reasonably accurate.  In the last few years, motivated by the search for gravitational waves, the knowledge of spin effects in the two-body dynamics and gravitational-wave emission within the post-Newtonian (PN)~\\footnote{We refer to $n$PN as the order equivalent to terms $O(c^{-2n})$ in the equations of motion beyond the Newtonian acceleration.} approximation has improved considerably. In particular, spin-orbit (SO) effects in the two-body equations of motion are currently known through 3.5PN order (i.e., 2PN order beyond the leading SO term)~\\cite{kidder93,OTO98,TOO01,Faye-Blanchet-Buonanno:2006, DJSspin,Porto:2010tr,Perrodin:2010dy,Levi:2010zu,Hartung:2011te}, and in the energy flux through 3PN order~\\cite{kidder95,OTO98,TOO01,W05,Blanchet-Buonanno-Faye:2006,Blanchet:2011zv} (i.e., 1.5PN order beyond the leading SO term). Moreover, spin-spin (SS) effects have been computed through 3PN order (i.e., 1PN order beyond the leading SS term) in the conservative dynamics~\\cite{kidder93,kidder95,MVGer05,Porto06,Porto:2006bt,Hergt:2007ha,SHS07, hergt_schafer_08,SHS08,SSH08,Porto:2008tb,PR08,Levi:2010} and also in the multipole moments~\\cite{Porto:2010zg}.  In order to build reliable templates and search for gravitational-waves from high-mass compact binaries that merge in the detector bandwidth, it is crucial to improve the PN approximation by resumming the dynamics and gravitational emission in a suitable way and by using numerical relativity and perturbation theory as a guidance. The effective-one-body approach (EOB)~\\cite{Buonanno00, Buonanno99, DJS00, Damour01c, Buonanno06} offers the possibility of fulfilling this goal. The EOB approach uses the results of PN theory, not in their original Taylor-expanded form (i.e., as polynomials in $v/c$), but instead in a suitably resummed form. In particular, it maps the dynamics of two compact objects of masses $m_1$ and $m_2$, and spins $\\mathbf{S}_1$ and $\\mathbf{S}_2$, into the dynamics of one test particle of mass $\\mu = m_1\\,m_2/(m_1+m_2)$ and spin $\\mathbf{S}_*$ moving in a deformed Kerr metric with mass $M = m_1+m_2$ and spin $\\mathbf{S}_\\mathrm{Kerr}$. The deformation parameter is the symmetric mass ratio $\\eta=m_1\\,m_2/(m_1+m_2)^2$, which ranges between $0$ (test particle limit) and $1/4$ (equal-mass limit). The analyses and theoretical progress made in Refs.~\\cite{Buonanno-Cook-Pretorius:2007, Buonanno2007, Pan2007, Boyle2008a, Buonanno:2009qa, Barausse:2009xi, Pan:2009wj, Pan2010hz, Damour2007a, DN2007b, DN2008, DIN, Damour2009a, Yunes:2009ef, Yunes:2010zj,Bernuzzi:2010xj,Pan:2011gk} have demonstrated that faithful EOB templates describing the full signal (i.e., the inspiral, merger and ringdown) can be built and used in real searches~\\cite{Abadie:2011kd}.  Here we build on previous work~\\cite{Damour01c,Damour:2007nc,Barausse:2009aa,Barausse:2009xi}, employ the recent results of Ref.~\\cite{Hartung:2011te} and extend the EOB conservative dynamics, i.e. the EOB Hamiltonian, through 3.5PN order in the SO couplings. Since the mapping between the PN-expanded Hamiltonian (or real Hamiltonian) and the EOB Hamiltonian is not unique, we explore two specific classes of EOB Hamiltonians, which differ by the way the spin variables of the real and effective descriptions are mapped.  This paper is organized as follows. In Sec.~\\ref{sec:SEOBH}, after reviewing the logic underpinning the construction of the EOB Hamiltonian, we proceed in steps and extend the EOB Hamiltonian proposed in Ref.~\\cite{Barausse:2009xi} through 3.5PN order in the SO couplings. In particular, in Sec.~\\ref{sec:ADMH} we derive the PN-expanded Arnowitt-Deser-Misner (ADM) Hamiltonian in the EOB canonical coordinates; then, after computing in Sec.~\\ref{sec:EffH} the effective Hamiltonian corresponding to the canonically transformed PN-expanded ADM Hamiltonian, we compare it (in Sec.~\\ref{sec:PNexp}) to the deformed-Kerr Hamiltonian for a spinning test-particle~\\cite{Barausse:2009xi}, and work out (in Secs.~\\ref{sec:SEOBgeo} and \\ref{sec:SEOBngeo}) two classes of EOB Hamiltonians. In Sec.~\\ref{sec:EOBdynamics} we study the dynamics of these Hamiltonians for equatorial orbits, and in Sec.~\\ref{sec:concl} we summarize our main conclusions.  We use geometric units $G = c = 1$ throughout the paper, except when performing PN expansions, where powers of the speed of light $c$ are restored and play the role of book-keeping parameters. ",
        "conclusions": "\\label{sec:concl}  Recently, Ref.~\\cite{Hartung:2011te} has computed the 3.5PN SO effects in the ADM Hamiltonian. We have taken advantage of this result and extended the EOB Hamiltonian of spinning black holes to include these higher-order SO couplings.  Building on previous work~\\cite{Damour01c,Damour:2007nc}, and in particular on the EOB Hamiltonian of Refs.~\\cite{Barausse:2009aa,Barausse:2009xi}, which reproduces the SO test-particle couplings exactly at all PN orders, we have worked out two classes of EOB Hamiltonians, which differ by the way the spin variables are mapped between the effective and real descriptions. One class of EOB Hamiltonians is the straightforward extension to the next PN order of the EOB Hamiltonian of Ref.~\\cite{Barausse:2009xi}. It uses a mapping between the real and effective spin variables that depends on the dynamical orbital variables $\\mathbf{p}^2$, $\\mathbf{n} \\cdot \\mathbf{p}$  and $r$. By contrast, the other class of EOB Hamiltonians uses a mapping between the real and effective spin variables that {\\it does not} depend on these dynamical orbital variables. We achieved this result at the cost of modifying the Hamilton-Jacobi equation of a spinning test-particle.  Quite interestingly, when restricting to spins aligned or antialigned with the orbital angular-momentum and to equatorial circular orbits, we find that the predictions of these two classes of EOB Hamiltonians for the ISCO frequency, energy and angular momentum, and for the maximum of the orbital frequency during the plunge are generally similar. However, for high spins the model with dynamical mapping of the spins may present somewhat lower maximum frequencies and larger ISCO frequencies.  As pointed out originally in Ref.~\\cite{Damour:2007nc}, several gauge parameters can enter the canonical transformation that maps the real and effective Hamiltonians. If the Hamiltonian were known exactly, i.e., at all PN orders, then physical effects should not depend on these parameters. However, since we know the Hamiltonian only at a certain PN order, we expect these gauge parameters to lead to non-negligible differences. In fact, we obtained that when setting all the gauge parameters to zero, the maximum frequency during the plunge has a non-monotonic dependence on the spins, and varies quite significantly as a consequence of the inclusion of the 3.5 PN SO couplings. We found instead that when choosing the gauge parameters so that the terms depending on the radial momentum disappear from our spin mapping (in the model with dynamical spin mapping) or from the modifications to the Hamilton-Jacobi equation (in the  model with non-dynamical spin mapping), the maximum frequency during the plunge has a much more regular behavior and varies by small amounts when adding the 3.5PN SO couplings. This suggests that such a choice of the gauge parameters may accelerate the convergence of the model's results in the strong-field region where the plunge takes place.  The EOB Hamiltonians derived in this paper can be calibrated to numerical-relativity simulations with the goal of building analytical templates for LIGO and Virgo searches. A first example was obtained in Ref.~\\cite{Pan:2009wj}, where the EOB Hamiltonian at 2.5PN order in the SO couplings of Ref.~\\cite{Damour:2007nc} was calibrated to two highly-accurate numerical simulations. Results that use the EOB Hamiltonian at 3.5PN order developed in this paper will be reported in the near future~\\cite{Taracchini:2011}.  Lastly, while finalizing this work,  Ref.~\\cite{2011arXiv1106.4349N} appeared in the archives as a preprint. Both this paper and Ref.~\\cite{2011arXiv1106.4349N} derive the effective gyromagnetic coefficients [see Eq.~(\\ref{gyroeffSO})], but with two different methods. Our computation uses the Lie method to generate both the purely-orbital and the spin-dependent canonical transformations, while  Ref.~\\cite{2011arXiv1106.4349N} first applies explicitly the  purely-orbital transformation from ADM to EOB coordinates, and then uses Eq.~\\eqref{Htransf} to account for the effect of a spin-dependent canonical transformation. As a result of these different procedures, and as discussed in Sec.~\\ref{sec:ADMH}, the 2.5PN gauge parameters in our spin-dependent canonical transformation coincide with those of Ref.~\\cite{2011arXiv1106.4349N}, but the 3.5PN gauge parameters have different meanings in the two approaches and therefore do not coincide. However, by suitably expressing our 3.5PN gauge parameters in terms of those of Ref.~\\cite{2011arXiv1106.4349N}, we find that our effective gyromagnetic coefficients fully agree with those of Ref.~\\cite{2011arXiv1106.4349N}. This amounts to saying that our gyromagnetic coefficients agree with those of Ref.~\\cite{2011arXiv1106.4349N} up to a canonical transformation, and are therefore physically equivalent. More importantly, in this paper we have focused on and worked out two classes of EOB Hamiltonians that are different from the one considered in Ref.~\\cite{2011arXiv1106.4349N}."
    },
    "1107/1107.1593_arXiv.txt": {
        "abstract": "{The search for a diffuse neutrino flux component from astrophysical sources complements searches for point sources. In 2010 and 2011 many new results have been published on this subject. Realistic models can now be tested with these new measurements from the leading neutrino telescope experiments. An overview over these recent results is given.}  \\normalsize\\baselineskip=15pt ",
        "introduction": "Neutrinos are produced in a variety of sources throughout the universe. The most prominent natural neutrino point source is our Sun. By exploring the angular correlation between the neutrino arrival direction and the position of the Sun, the solar origin of these neutrinos could be undoubtedly proven. However apart from our Sun no other permanent neutrino sources could be identified so far. An alternative is the search for a diffuse neutrino flux. Individual sources cannot be distinguished but the spectral shape of their integral contribution might be used to distinguish them from possible background. Such a diffuse neutrino flux might exist at very different energies.  A large contribution is predicted from neutrinos which decoupled from thermal equilibrium at a temperature of the universe of about 1~MeV. Due to the expansion of the universe these neutrinos (equipartitioned over all flavours) are expected to have today a black body radiation spectrum with a mean energy in the meV region. From measurements of neutrino oscillations it is known that the two heavier neutrinos have masses of $m_2>9$~meV$/c^2$ and $m_3>50$~meV$/c^2$. Therefore they must be non-relativistic today. So far no experimental technique is known to detect them.  Another diffuse neutrino flux must exist when summing the neutrino yield of all stellar fusion processes in the history of the universe. These neutrinos are expected to have spectra similar to solar neutrinos with energies mainly in the region of 0.1~MeV to 10~MeV. At slightly higher energies (up to 50~MeV) one should find a diffuse flux from neutrinos which had been produced during neutron star formation at core collapse Supernovae explosions throughout the history of the universe. A search for such a  neutrino flux has been done in the SuperKamiokande experiment~\\cite{SuperK-diff}. No signal is observed and a flux limit could be set. This is discussed in more detail elsewhere in these proceedings~\\cite{SuperK}.  In the following we concentrate on the diffuse flux at energies $E_\\nu > 1$~TeV. Only the most violent processes in the universe such as Active Galactic Nuclei (AGNs) or Gamma Ray Bursts (GRBs) can contribute in this energy range. Its observation requires significantly larger detectors compared to those found in underground sites. In the following, results from large scale Neutrino Telescopes are presented which use natural transparent media such as deep sea water or Antarctic ice. A detailed introduction into the concept and physics of these devices is found elsewhere in these proceedings~\\cite{NT}. Here results from the Baikal~\\cite{Baikal-diff-casc}, AMANDA~\\cite{Amanda-diff-mu}, ANTARES~\\cite{Antares} and IceCube~\\cite{IceCube} Neutrino Telescopes will be discussed. ",
        "conclusions": "During the last year a wealth of new results have been published on the search for diffuse neutrino fluxes from astrophysical sources. These searches complement point source searches. So far no excess of events beyond the expected rate of atmospheric neutrino events has been found. The sensitivity of the analyses is already high enough to constrain or exclude various models. For the first time fluxes below the W\\&B upper bound~\\cite{WB-limit} are tested. Fast publication has become a standard which illustrates a good understanding of the used detectors and media."
    },
    "1107/1107.3844_arXiv.txt": {
        "abstract": "{The migration of planets plays an important role in the early planet-formation process. An important problem has been that standard migration theories predict very rapid inward migration, which poses problems for population synthesis models. However, it has been shown recently that low-mass planets (20-30 $M_{Earth}$) that are still embedded in the protoplanetary disc can migrate outwards under certain conditions. Simulations have been performed mostly for planets at given radii for a particular disc model. } {Here, we plan to extend previous work and consider different masses of the disc to quantify the influence of the physical disc conditions on planetary migration. The migration behaviour of the planets will be analysed for a variety of positions in the disc. } {We perform three-dimensional (3D) radiation hydrodynamical simulations of embedded planets in protoplanetary discs. We use the explicit-implicit 3D hydrodynamical code {{\\tt NIRVANA}} that includes full tensor viscosity, and implicit radiation transport. For planets on circular orbits at various locations we measure the radial dependence of the torques for three different planetary masses. } {For all considered planet masses (20-30 $M_{Earth}$) in this study we find outward migration within a limited radial range of the disc, typically from about 0.5 up to 1.5-2.5 $a_{Jup}$. Inside and outside this interval, migration is inward and given by the Lindblad value for large radii. Interestingly, the fastest outward migration occurs at a radius of about $a_{Jup}$ for different disc and planet masses. Because outward migration stops at a certain location in the disc, there exists a zero-torque distance for planetary embryos, where they can continue to grow without moving too fast. For higher disc masses ($M_{disc} > 0.02 M_\\odot$) convection ensues, which changes the structure of the disc and therefore the torque on the planet as well. } { Outward migration stops at different points in the disc for different planetary masses, resulting in a quite extended region where the formation of larger cores might be easier. In higher mass discs, convection changes the disc's structure resulting in fluctuations in the surface density, which influence the torque acting on the planet, and therefore its migration rate. Because convection is a 3D effect, 2D simulations of massive discs with embedded planets should be handled with care. } ",
        "introduction": "\\label{sec:introduction} Planetary migration of embedded low-mass planets in isothermal discs indicates inward migration, so that the planet might be lost in the star before the accretion disc is gone \\citep{2002ApJ...565.1257T}. Recent studies (starting with the work of \\citet{2006A&A...459L..17P}) have shown that the inclusion of radiation transport in planet-disc interaction studies resulted in a slowed down or even reversed migration \\citep{2008ApJ...672.1054B, 2008A&A...485..877P, 2008A&A...478..245P, 2008A&A...487L...9K, 2009A&A...506..971K, 2010MNRAS.408..876A, 2011arXiv1103.3502A}.  All authors agree that the inclusion of radiation transport is an important effect that should be considered, however, not all authors find outward migration. The efficiency of outward migration  depends on the magnitude of positive corotation torques. These are determined by the local entropy and vortensity gradients. For isothermal and adiabatic discs these gradients cannot be maintained and so-called torque saturation reduces the corotation effects. However, the inclusion of radiative transport and viscosity prevents saturation, and the negative Lindblad torques (caused by the spiral arms) can be overwhelmed by the positive contributions from the corotation region.   The magnitude of the effect depends on the planetary mass \\citep{2008A&A...487L...9K, 2009A&A...506..971K} and the temperature gradient in the disc. For example, \\citet{2011arXiv1103.3502A} found for temperature slopes of $\\beta > 1.0$ (with $T \\propto r^{-\\beta}$) that outward migration is possible even for p lanets with up to $50 M_{Earth}$ in 3D simulations. The migration rate increases with an increasing temperature slope. These results are also reflected by the theoretical analysis from \\citet{2010ApJ...723.1393M}. Recently, improved theoretical torque formulae for low-mass  planets embedded in an adiabatic disc have been presented by \\citet{2009ApJ...703..857M} and \\citet{2010MNRAS.401.1950P}. The most recent improvements consider the necessary inclusions of radiative diffusion and viscosity \\citep{2010ApJ...723.1393M, 2011MNRAS.410..293P}. These have been developed for 2D discs where the diffusive effects can operate only in the disc's plane. Interestingly, radiative cooling alone can lead to similar, even stronger effects \\citep{2008A&A...487L...9K}.  A very important open question is how far the planet will move outwards in a fully radiative disc. When the protoplanets are stopped from migrating outwards at a certain point in the disc, a protoplanet moving inwards from farther out will stop at the same location in the disc. Therefore, a planetary trap inside the disc can be created, which can act as an area of planetary mergers, leading to larger cores. A planetary trap inside the disc created by surface density changes can function as a feeding zone to planetary cores \\citep{2008A&A...478..929M}, but it is unclear how realistic these surface density changes are. The inclusion of radiation transport/cooling in a disc might provide such a trap in a normal disc structure, at least for planets within a given mass range.  In \\citet{2011ApJ...728L...9S} the authors show in N-body simulations that because of the inclusion of unsaturated type-I-migration \\citep{2010MNRAS.401.1950P} large planetary cores (up to $10 M_{Earth}$) can form in protoplanetary discs well before the disc is accreted onto the star.  In population synthesis models that include the standard migration prescription for isothermal discs the migration is too fast, such that the outcome distribution does not match the observed one. Specifically, the type-I-migration rate should be $10$ to $1000$ times slower than expected from linear analysis \\citep{2004A&A...417L..25A, 2008ApJ...673..487I, 2009A&A...501.1139M}. As pointed out above, this could be provided by the inclusion of radiation transport in hydrodynamical type-I-migration simulations. An updated version of the type-I-migration, by using the formula from \\citet{2010MNRAS.401.1950P}, in population synthesis models shows very promising results \\citep{2011arXiv1101.3238M}, but additional simulations are needed to verify these assumptions.  In our previous simulations and studies \\citep{2009A&A...506..971K, 2010A&A.523...A30, 2011A&ABitschKley} we have only considered one standard disc model with a given mass. In reality, of course, protoplanetary discs can have a variety of masses. A fully radiative disc will settle in an equilibrium state between viscous heating and radiative transport/cooling. Given that all other physical parameters are fixed, the resulting disc structure only depends on the disc's mass and viscosity, because a more massive or more viscous disc is able to heat the disc more effectively. A more massive disc therefore influences the migration rate of embedded planets. Here we focus on low-mass planets that can undergo outward migration in fully radiative discs.  In this paper we extend previous studies and investigate the possible radial range over which outward migration can occur and analyse the influence of the disc's mass on the migration in detail. An important effect is the onset of convection in the disc, which becomes stronger for more massive discs.  In Section \\ref{sec:setup} we give a short overview of our code and numerical methods, where more details can be found in \\citep{2009A&A...506..971K}. The influences of the distance of the embedded planet to the central star is discussed in Section \\ref{sec:outwardrange}. We then analyse the influence of the disc's mass on the disc's structure (density, temperature, aspect ratio) and then the influence on migration of embedded low-mass planets in Section \\ref{sec:discmass}. Convection in the disc is also briefly discussed in Section \\ref{sec:discmass}. We then summarize and conclude in Section \\ref{sec:Sumcon}. ",
        "conclusions": "\\label{sec:Sumcon}  We performed full 3D radiation hydrodynamical simulations of low-mass planets embedded in accretion discs at different distances to the central star and for various disc masses.  In the first sequence of our simulations we changed the planetary distance to the central star of embedded planets on circular orbits. With increasing distance to the central star, the torque acting on $20 M_{Earth}$ planets embedded in fully radiative discs becomes even more reduced and it reaches negative torques for longer distances. We find an equilibrium, zero-torque distance, to the central star for $20 M_{Earth}$ planets at $r \\approx 2.4 a_{Jup}$. This equilibrium distance varies with the planetary mass (for $25 M_{Earth}$ planets it is $r \\approx 1.9 a_{Jup}$ and $r \\approx 1.4 a_{Jup}$ for $30 M_{Earth}$ planets), indicating that a quite extended region in the disc might act as a feeding zone to create even larger planetary cores. The concept of equilibrium radius (zero torque radius) for planetary embryos in fully unsaturated discs has been stated in \\citet{2010ApJ...715:L68} as well and it could easily act as a feeding or collection zone for planetary embryos.  Planets embedded in fully radiative discs migrating outwards create a very sensitive pattern in the surface density distribution. Ahead and inside of the planet a density increase is visible (see second panel from the top in Fig.~\\ref{fig:Mig2DRho}), which shrinks with increasing distance to the central star. This density enhancement is indeed accountable for the positive torque acting on the planet (also visible as a spike in the radial torque density distribution in Fig.~\\ref{fig:MigGamma3D}), but as the distance to the central star increases, this effect becomes less, so that it cannot overcompensate the negative Lindblad torques any more, which results in inward migration.  We compared our results to the recently developed torque formulae by \\citet{2010MNRAS.401.1950P,2011MNRAS.410..293P} and \\citet{2010ApJ...723.1393M}. \\citet{2010MNRAS.401.1950P} includes just the fully unsaturated torques in the inviscid and adiabatic case, shows no torque reversal option for constant $\\beta$ and is as such unphysical when comparing it with the long-term evolution of planets. However, as in our simulations $\\beta$ changes in the disc with distance to the central star ($\\beta (r=1.0 a_{Jup}) =1.7$ to $\\beta (r=5.0 a_{Jup}) = 0.33$), the torque given by the formula becomes negative for large distances to the central star. But, overall the match of the torque of the formula with our simulations is not good. This formula is only valid in the first orbits after the planet is embedded when the torques are still unsaturated, however, the torques do saturate in time.  However, as expected, the improved version of \\citet{2011MNRAS.410..293P} that includes viscous and heat diffusion describes our results more accurately. It also features a transition from positive to negative torques, but the negative torques at large distances to the central star are about a factor of two larger than in our simulations. Even though an exact match of the torques has not been achieved, the trend seems to be caputured quite well by the formula.  The torque formulae are derived for 2D discs with thermal diffusion operating only in the disk's midplane, while our simulations are 3D, and take full account of vertical diffusion as well, which can change the structure of the disc. The formulae were also derived for given gradients in temperature and surface density, but in a real disc the temperature and surface density profiles are disturbed when a planet is embedded in a disc. Because the formulae were derived and checked for a $5 M_{Earth}$ planet (with about $20$ per cent agreement), the disturbances of such a small planet in a disc are much weaker than for our embedded $20 M_{Earth}$ planet, which may give rise to more significant, non-linear disturbances in the temperature and density profiles. All of this may lead to differences between our simulations and the theoretical formulae in the torque acting on the planet.   The formula of \\citet{2010ApJ...723.1393M} does not match our simulations that well. The torques from the formula are always negative, which is in contrast to our simulations, where we find outward migration for $r<2.5 a_{Jup}$. For long distances to the central star, \\citet{2010ApJ...723.1393M} match better than \\citet{2011MNRAS.410..293P}, as the Lindblad torque is reduced. The Lindblad torque of \\citet{2010ApJ...723.1393M} also matches quite well with the isothermal torque of \\citet{2010ApJ...724..730D}, when taking the factor $\\gamma$ into account due to the different sound speeds in isothermal and fully radiative discs. The overall trend to have torques from positive to negative seems best achieved with \\citet{2011MNRAS.410..293P}, but a quantitative agreement is still lacking. In Appendix \\ref{app:comp} we discuss the influence of the smoothing length on the formulae.    For increasing disc masses, the temperature, density (in midplane), and aspect ratio of the disc increases in the equilibrium state where  viscous heating and radiative transport/cooling are in balance. The convective zone in the inner discs stretches farther out with increasing disc mass, resulting in high fluctuations of the surface density in our computed domain for discs with a mass higher than $M_{disc} \\approx 0.02 M_\\odot$.  Starting from a $M_{disc}=0.01 M_\\odot$ disc, the torque acting on embedded $20 M_{Earth}$ planets decreases for increasing disc masses. As the disc mass increases, the convective zone in the disc stretches farther out from the central star and influences planetary migration. The fluctuations in the disc's density disrupt the torque acting on the planet on a stationary orbit for high-mass discs in a way that the torque is very irregular and shows high fluctuations as well, making it difficult to determine the correct direction of migration. For lower disc masses, the torque reduces as well, presumably because of the same reasons as the torque is reduced for longer distances to the central star in a discs with $M_{dics} = 0.01 M_\\odot$.  The formula in \\citet{2011MNRAS.410..293P} fits within a factor of three with our 3D simulations for planets in discs with $M_{dics} \\approx 0.01 M_\\odot$. For higher disc masses, the difference between the formula and our 3D simulations becomes smaller. However, the 3D simulations show a decreasing torque for increasing disc mass, while the \\citet{2011MNRAS.410..293P} formula increases slightly. As the disc mass increases, viscous heating increases and cooling becomes inefficient, which results in a structure similar to an adiabatic disc. In adiabatic discs the corotation torques saturates, resulting in a lower torque acting on the planet, hence the drop of the torque for increasing disc masses. Convection certainly plays a role in more massive discs, but it is unaccounted for in \\citet{2011MNRAS.410..293P}. In more massive discs the convective zone reaches longer distances from the central star, disrupting the density pattern near the embedded planet and thus creating fluctuations in the torque acting on the planet. These disruptions in the density pattern are caused by the convective cells evolving in the disc. These cells also change in time, giving rise to the stronger fluctuations of the torque.    Convection is inefficient for transporting angular momentum \\citep{1993ApJ...416..679K,2010MNRAS.404..L64}, but the influences of convection on planetary migration are very dramatic, because a planet close to the star in the convective zone of the disc is essentially disrupted. The direction of migration is not clearly determinable any more, but when the planet is farther out in the massive disc and the convection fades away, the direction of migration is easy to specify, indicating outward migration. Therefore the zero-torque radius for migration lies farther out in more massive discs.  Convection is also a 3D effect only and cannot be simulated in 2D discs (in $r$-$\\phi$ direction in the midplane). Two-dimensional simulations of planets in massive discs ($M_{disc} > 0.02 M_\\odot$) might therefore be inaccurate near the central star, because the effects of convection are not considered.  \\appendix"
    },
    "1107/1107.0656_arXiv.txt": {
        "abstract": "{Active galactic nuclei show a wealth of interesting physical processes, some of which are poorly understood. In a broader context, they play an important role in processes that are far beyond their immediate surroundings, owing to the high emitted power.} {We want to address a number of open questions, including the location and physics of the outflow from AGN, the nature of the continuum emission, the geometry and physical state of the X-ray broad emission line region, the Fe-K line complex, the metal abundances of the nucleus and finally the interstellar medium of our own Galaxy as seen through the signatures it imprints on the X-ray and UV spectra of AGN.} {We study one of the best targets for these aims, the Seyfert 1 galaxy Mrk~509 with a multiwavelength campaign using five satellites (XMM-Newton, INTEGRAL, Chandra, HST and Swift) and two ground-based facilities (WHT and PAIRITEL). Our observations cover more than five decades in frequency, from 2~$\\mu$m to 200~keV. The combination of high-resolution spectroscopy and time variability allows us to disentangle and study the different components. Our campaign covers 100 days from September to December 2009, and is centred on a simultaneous set of deep XMM-Newton and INTEGRAL observations with regular time intervals, spanning seven weeks.} {We obtain a continuous light curve in the X-ray and UV band, showing a strong, up to 60\\% flux increase in the soft X-ray band during the three weeks in the middle of our deepest monitoring campaign, and which is correlated with an enhancement of the UV flux. This allows us to study the time evolution of the continuum and the outflow. By stacking the observations, we have also obtained one of the best X-ray and UV spectra of a Seyfert galaxy ever obtained. In this paper we also study the effects of the spectral energy distribution (SED) that we  obtained on the photo-ionisation equilibrium. Thanks to our broad-band coverage, uncertainties on the SED do not strongly affect the determination of this equilibrium.} {Here we present our very successful campaign and in a series of subsequent papers we will elaborate on different aspects of our study.} ",
        "introduction": "Active galactic nuclei (AGN) are compact sources with very high luminosities, located at the centres of galaxies. Accretion onto the super-massive black holes (SMBH) at their centres is generally believed to be the driving process for the activity. Thanks to their brightness, they form one of the richest laboratories for studying astrophysical processes. In this paper we present one of the deepest multiwavelength campaigns of an AGN, the Seyfert 1 galaxy \\object{Mrk~509}. It is among the best for these studies because it is unique in combining X-ray brightness, outflow features, and significant but moderate variability. Below we introduce the most important astrophysical processes that are addressed by our study.  After our introduction of the relevant astrophysics, we briefly provide an overview in Sect.~\\ref{sect:mrk509} of the target of our campaign, Mrk~509, followed by an overview of the observations in Sect.~\\ref{sect:log}, and we present the spectral energy distribution (SED) and light curve in the subsequent sections. Details about various aspects of our study are deferred to subsequent papers of this series. ",
        "conclusions": "In this paper we have introduced our multiwavelength study of the Seyfert 1 galaxy Mrk~509, which used XMM-Newton, INTEGRAL, Chandra, HST, Swift, WHT and PAIRITEL. Our observations spanned 100 days of monitoring from September 2009 to December, covering band-passes from 2~$\\mu$m to 200~keV. The core of our programme consisted of ten simultaneous observations with XMM-Newton and INTEGRAL, followed by a long Chandra LETGS spectrum obtained simultaneously with an HST/COS far-UV spectrum. The high-resolution spectroscopy combined with the time variability monitoring enabled us to disentangle the different absorption and emission components in Mrk~509.  Using our data we produced a continuous light curve in the X-ray and UV bands and a comprehensive SED. Thanks to our broad-band coverage, we deduced that uncertainties on the SED have little effect on the photo-ionisation equilibrium that applies to our subsequent models of the ionised outflow. Our light curve shows a 60\\% flux increase in the soft X-ray band correlated with an enhancement of the UV flux.  A series of subsequent papers will elaborate on these results and others from our campaign. The stacked XMM-Newton RGS spectrum is presented in \\citet[][paper~II]{kaastra2011}. The time-averaged ionisation and velocity structure of the outflow deduced from this spectrum is presented by \\citet[][paper~III]{detmers2011}, the broad-band continuum variability by \\citet[][paper~IV]{mehdipour2011}. \\citet[][paper~V]{ebrero2011} describes the Chandra LETGS data and \\citet[][paper~VI]{kriss2011} presents the analysis of the HST/COS data. More papers are in preparation."
    },
    "1107/1107.2653_arXiv.txt": {
        "abstract": "We take advantage of the sensitivity and resolution of the {\\em Herschel Space Observatory} at $100$ and $160$\\,$\\mu$m to directly image the thermal dust emission and investigate the infrared luminosities ($\\lir$) and dust obscuration of typical star-forming ($L^{\\ast}$) galaxies at high redshift.  Our sample consists of $146$ UV-selected galaxies with spectroscopic redshifts $1.5\\le z_{\\rm spec}<2.6$ in the GOODS-North field.  Supplemented with deep Very Large Array (VLA) and {\\em Spitzer} imaging, we construct median stacks at the positions of these galaxies at $24$, $100$, and $160$\\,$\\mu$m, and $1.4$\\,GHz.  The comparison between these stacked fluxes and a variety of dust templates and calibrations implies that typical star-forming galaxies with UV luminosities $\\luv \\ga 10^{10}$\\,L$_{\\odot}$ at $z\\sim 2$ are luminous infrared galaxies (LIRGs) with a median $\\lir= (2.2\\pm 0.3)\\times 10^{11}$\\,L$_{\\odot}$. Their median ratio of $\\lir$ to rest-frame $8$\\,$\\mu$m luminosity ($\\mirlum$) is $\\lir/\\mirlum = 8.9\\pm1.3$ and is $\\approx 80\\%$ larger than that found for most star-forming galaxies at $z\\la 2$.  This apparent redshift evolution in the $\\lir/\\mirlum$ ratio may be tied to the trend of larger infrared luminosity surface density for $z\\ga 2$ galaxies relative to those at lower redshift.  Typical galaxies at $1.5\\le z<2.6$ have a median dust obscuration $\\lir/\\luv = 7.1\\pm1.1$, which corresponds to a dust correction factor, required to recover the bolometric star formation rate (SFR) from the unobscured UV SFR, of $5.2\\pm0.6$.  This result is similar to that inferred from previous investigations of the UV, H$\\alpha$, $24$\\,$\\mu$m, radio, and X-ray properties of the same galaxies studied here.  Stacking in bins of UV slope ($\\beta$) implies that $L^{\\ast}$ galaxies with redder spectral slopes are also dustier, and that the correlation between $\\beta$ and dustiness is similar to that found for local starburst galaxies. Hence, the rest-frame $\\simeq 30$ and $50$\\,$\\mu$m fluxes validate on average the use of the local UV attenuation curve to recover the dust attenuation of typical star-forming galaxies at high redshift.  In the simplest interpretation, the agreement between the local and high redshift UV attenuation curves suggests a similarity in the dust production and stellar and dust geometries of starburst galaxies over the last $10$\\,billion years. ",
        "introduction": "\\label{sec:intro}  A key aspect of measuring total star formation rates is to assess how a galaxy's spectral energy distribution is modulated by the effects of interstellar dust.  The ability to quantify dust attenuation at high redshift is hindered by the limited dynamic range in luminosity at wavelengths that are sensitive to dust emission.  While the Lyman Break technique \\citep{steidel95} has proven to be the most effective method of identifying galaxies across a large dynamic range in luminosity and lookback time, it requires observations in the rest-frame UV: the massive stars giving rise to the UV continuum also produce much of the dust that attenuates this continuum.  The diminution of UV light is compounded by the fact that much of this radiation is re-emitted in the far-infrared where current instrumental sensitivity is insufficient to directly detect typical star-forming galaxies at $z\\ga 2$.  It has therefore become common to use local relations between monochromatic and bolometric luminosity to infer the star formation rates and dust attenuation of high-redshift galaxies.  Prior to the era of large-scale multi-wavelength surveys like GOODS \\citep{dickinsongoods03,giavalisco04}, inferring the dust attenuation at high redshift commonly entailed using relations between the variation of the UV continuum slope with dustiness as found for local starburst galaxies \\citep{meurer99, calzetti00, heckman98}.  However, such correlations were previously untested at high redshift.  The correlation between the redness of the UV slope ($\\beta$) and dust attenuation has limited applicability when examined over a larger range in galaxy stellar population and luminosity.  Galaxies with older and less massive stars contributing significantly to the UV emission can also exhibit a redder spectral slope \\citep{calzetti97, kong04, buat05}.  Further, the UV slope decouples from extinction for very luminous galaxies where virtually all of the star formation is obscured, such as is the case for low redshift ultra-luminous infrared galaxies (ULIRGs; \\citealt{goldader02}).  Nonetheless, the local trend between UV slope and dustiness is still applied widely to galaxies at high redshift; this correlation is often the only means by which one can infer the dust attenuation of galaxies at $z\\ga 3$, where the dust emission from a typical galaxy may be several orders of magnitude below the sensitivity limits of current infrared instruments.  Incremental progress in determining dust attenuation at high redshift was achieved with the first ultra-deep radio and X-ray surveys (e.g., \\citealt{richards00, alexander03}).  Though the sensitivity (even at GOODS-depth) at these wavelengths was insufficient to individually detect typical star-forming galaxies at $z\\ga 2$, these surveys did allow for estimates of the mean dust attenuation based on stacking analyses.  Initial studies suggested a general agreement between UV and radio/X-ray inferences of dust attenuation at $z\\sim 2-3$ \\citep{nandra02, seibert02, reddy04, reddy05a, reddy06a, daddi07a, pannella09}.  The launch of the {\\em Spitzer Space Telescope} enabled the first direct detection of the dust emission in non-lensed $L^{\\ast}$ galaxies at $z\\sim 2$; at these redshifts, {\\em Spitzer}'s MIPS $24$\\,$\\mu$m band is sensitive to the strongest dust emission feature (at $7.7$\\,$\\mu$m) in star-forming galaxies.  This feature arises from the stochastic UV photo-heating of small dust grains and hydrocarbons (e.g., \\citealt{puget89, tielens99}) and is found to correlate with the UV radiation from OB stars (e.g., \\citealt{forster04, roussel01}), albeit with significant scatter (e.g., \\citealt{kennicutt09, hogg05, helou01, engelbracht05, normand95}).  Several studies have demonstrated that the dust attenuation inferred from the UV spectral slope, $\\beta$, for typical $z\\sim 2$ galaxies is in general agreement with that inferred from MIPS $24$\\,$\\mu$m (e.g., \\citealt{reddy06a, reddy10a, daddi07a}).  Regardless, local studies of resolved galaxies have emphasized the complexity of the $8$\\,$\\mu$m emission and the need to combine it with other obscured (IR) and unobscured (UV, H$\\alpha$) tracers of star formation in order to obtain the most reliable measure of the bolometric luminosity \\citep{kennicutt09}.  To this end, we take advantage of the improved sensitivity of {\\em Herschel}/PACS at $100$ and $160$\\,$\\mu$m \\citep{pilbratt10, poglitsch10} to measure directly for the first time the thermal dust emission from a large sample of typical ($L^{\\ast}$) star-forming galaxies at $z\\sim 2$.  We make use of the deep PACS data made possible by the GOODS-{\\em Herschel} Open Time Key project (PI: D. Elbaz).  We supplement these data with deep Very Large Array (VLA) $1.4$\\,GHz continuum imaging in the GOODS-North field \\citep{morrison10}.  The primary aim is to constrain the average infrared luminosities of galaxies at high redshift, particularly those selected by their rest-frame UV colors, and to compare them to UV-based inferences.  We begin in Section~\\ref{sec:sample} by discussing the UV-selected galaxies and previous efforts to measure their stellar populations and dust attenuation.  We also provide a brief description of the {\\em Herschel} and radio data.  Our stacking method and stacking simulations are described in Section~\\ref{sec:stacking}, followed by a discussion of the dust spectral energy distribution (SED) fits to the stacked fluxes and the total infrared luminosities in Section~\\ref{sec:luminosities}.  In Section~\\ref{sec:discussion} we proceed to compare the {\\em Herschel}-based inferences of the dust attenuation with that inferred from the UV slope and discuss systematics in the dust obscuration with stellar population age and bolometric and UV luminosity.  For ease of comparison with other studies, we assume a \\citet{salpeter55} initial mass function (IMF) and adopt a cosmology with $H_{0}=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{\\rm m}=0.3$. ",
        "conclusions": "We have used the deep $100$ and $160$\\,$\\mu$m data of the GOODS-{\\em Herschel} Open Time Key Program, supplemented with deep {\\em Spitzer}/MIPS $24$\\,$\\mu$m and VLA 1.4\\,GHz radio imaging, to investigate the infrared luminosities and dust obscuration of typical star-forming galaxies at high redshift.  We focus on the median stacked mid-infrared, far-infrared, and radio fluxes of a sample of $146$ UV-selected galaxies with spectroscopic redshifts $1.5\\le z_{\\rm spec}<2.6$ in the GOODS-North field.  These galaxies have luminosities around $L^{\\ast}$ of the UV luminosity function at these redshifts.  Because these galaxies are individually undetected at $100$, $160$\\,$\\mu$m, and $1.4$\\,GHz, we perform median stacking analyses to measure their average fluxes.  Three tests are performed to verify the robustness of our stacked results.  The first test measured the effects of clustering. The second test determined the chance probability of measuring a stacked flux as high as the one obtained for the target galaxies.  The third test allows us to measure the biases and uncertainties in stacked flux measurements by stacking on the positions of artificial sources added to the images.  To interprete these fluxes, we consider a variety of dust templates, including those of \\citet{elbaz11}, \\citet{chary01}, \\citet{dale02}, and \\citet{rieke09}, as well as the $24$\\,$\\mu$m-$\\lir$ and radio-infrared correlations of \\citet{reddy10a} and \\citet{bell03}, respectively.  Fitting these templates to the bias-corrected fluxes and infrared colors reveals that $L^{\\ast}_{\\rm UV}$ galaxies at $z\\sim 2$ with UV luminosities $\\luv \\ga 10^{10}$\\,L$_{\\odot}$ have a median infrared luminosity of $\\lir= (2.2\\pm 0.3)\\times 10^{11}$\\,L$_{\\odot}$.  Galaxies in our sample exhibit $\\lir/\\mirlum$ (IR8) ratios that are a factor of $\\approx 2$ larger than those found for most star-forming galaxies at lower redshifts, likely due to the fact these UV-selected galaxies are relatively compact for their infrared luminosity.  One possibility is that the roughly constant IR8 ratio observed for most (main sequence) galaxies with redshifts $z\\la 2.0$ likely shifts towards larger values at $z\\ga 2.0$, due to the fact that the galaxies are on average smaller for their IR luminosity relative to lower redshift galaxies.  Stacking the {\\em Spitzer}, {\\em Herschel}, and VLA data as a function of UV spectral slope, $\\beta$, and bolometric luminosity, $\\lbol$, indicates that galaxies with redder $\\beta$ and higher $\\lbol$ have on average larger infrared luminosities, in accord with expectations based on trends found for local star-forming galaxies. Based on the $\\lir$, we proceed to examine the dust attenuation, $\\lir/\\luv$, and the dust correction factor (i.e., the factor required to recover the bolometric SFR from the unobscured UV SFR), 1+ SFR$_{\\rm IR}$/SFR$_{\\rm UV}$, for galaxies in our sample.  For $L^{\\ast}$ galaxies at $z\\sim 2$, we find a dust attenuation of $\\lir/\\luv = 7.1\\pm1.1$, which corresponds to a dust correction factor of $5.2\\pm 0.6$, implying that $\\simeq 80\\%$ of the star formation is obscured.  This result is consistent with those found in previous studies of the dust-corrected UV and H$\\alpha$, $24$\\,$\\mu$m, and radio and X-ray stacks of the same galaxies examined here \\citep{reddy04, reddy06a, reddy10a}.  We have examined the relationship between the UV spectral slope, $\\beta$, and dustiness for $z\\sim 2$ galaxies.  A small fraction ($\\approx 14\\%$) of galaxies are identified as having stellar population ages $\\la 100$\\,Myr from fitting stellar population SEDs to the observed optical through near-IR photometry.  The upper limit in dust attenuation for these ``young'' galaxies suggests that they may follow a ``steeper'' attenuation curve than the one observed for local starburst galaxies \\citep{meurer99, calzetti00}, in the sense that they are less dusty (i.e., have lower $\\lir/\\luv$ ratios) for a given $\\beta$ than older and more typical galaxies at the same redshift. This result was first noted in the $24$\\,$\\mu$m analyses of \\citet{reddy06a} and further investigated in \\citet{reddy10a}, and may be attributable to the higher dust covering fractions in young galaxies.  When considering more typical galaxies with ages $\\ga 100$\\,Myr, we find that their median $\\lir/\\luv$ increases as $\\beta$ becomes redder, and that this correlation is essentially identical to that found for local starbursts \\citep{meurer99, calzetti00}. Comparison of the {\\em Herschel} results at $z\\sim 2$ with measurements of local galaxies confirms the previously found trends that imply that LIRGs at $z\\sim 2$ are more UV transparent (i.e., less dusty) than LIRGs in the local universe, and that UV-faint galaxies at $z\\sim 2$ have lower $\\lir$ and hence fainter bolometric luminosities than UV-bright galaxies at $z\\sim 2$ (e.g., see also \\citealt{reddy10a}).  Based on the direct {\\em Herschel} measurements of the rest-frame $\\simeq 30$ and $50$\\,$\\mu$m thermal dust emission, we find that the local UV attenuation curve holds at $z\\sim 2$ for galaxies with $\\lbol \\la 10^{12}$\\,L$_{\\odot}$, suggesting a remarkable similarity in the processes governing star formation and dust production over the last $10$\\,billion years of cosmic time.  We have made significant progress in evaluating the effects of dust in typical star-forming galaxies at $z\\sim 2$ as traced by thermal infrared emission.  However, even the superb sensitivity and resolution of {\\em Herschel} are not sufficient to individually detect $L^{\\ast}$ galaxies at these redshifts.  The most accurate estimate of dust luminosity, and hence total star formation rate, can only come from combining {\\em individual} measurements of UV and IR luminosities.  Though the observations will be targeted, ALMA promises to extend our knowledge of the dust continuum luminosities of {\\em individual} star-forming galaxies at $z\\sim 2$, thus providing the key ingredient to combine with UV measurements.  These developments point to a wealth of forthcoming information on the properties of dust in typical star-forming galaxies at high redshift.  Ultimately, a robust characterization of the relationship between infrared and short wavelength (UV, H$\\alpha$) emission will be crucial for inferring dust attenuation and total star formation rates of the very faint and/or very high redshift galaxies that will be inaccessible to even the next generation of long wavelength observatories."
    },
    "1107/1107.2715_arXiv.txt": {
        "abstract": "Stereoscopic visualization is seldom used in Astrophysical publications and presentations compared to other scientific fields, e.g., Biochemistry, where it has been recognized as a valuable tool for decades. We put forth the view that stereo pairs can be a useful tool for the Astrophysics community in communicating a truer representation of astrophysical data. Here, we review the main theoretical aspects of stereoscopy, and present a tutorial to easily create stereo pairs using \\textsc{Python}. We then describe how stereo pairs provide a way to incorporate 3D data in 2D publications of standard journals. We illustrate the use of stereo pairs with one conceptual and two Astrophysical science examples: an integral field spectroscopy study of a supernova remnant, and numerical simulations of a relativistic AGN jet. We also use these examples to make the case that stereo pairs are not merely an ostentatious way to present data, but an enhancement in the communication of scientific results in publications because they provide the reader with a realistic view of multi-dimensional data, be it of observational or theoretical nature. In recognition of the ongoing 3D expansion in the commercial sector, we advocate an increased use of stereo pairs in Astrophysics publications and presentations as a first step towards new interactive and multi-dimensional publication methods. ",
        "introduction": "\\label{Sec:intro} \\emph{Stereoscopy} consists in giving a depth perception out of 2D material to the viewer, and the concept behind it is fairly simple: it requires sending distinct and carefully chosen images to each eye, without one eye noticing the images intended for the other. The notion of depth perception, or \\emph{stereopsis}, has been discussed as early as A.D. 280 by Euclid \\citep{Okoshi76}. While there has been some experimentation using sketching techniques before 1800 \\citep{Norling53}, the invention of photography by Niepce \\citep{Smith1877,Perrier34} in the beginning of the 19$^{\\mathrm{th}}$ Century marked the real start of extensive experimentation with stereoscopy. \\cite{Wheatstone1838} assembled one of the earliest known stereoscopes, but Brewster is usually attributed the construction of the first practical viewing device, now referred to as the \\emph{Brewster stereoscope} \\citep{Norling53}.  As predicted by \\cite{Scripture1899}, stereoscopy encountered quite a strong success in these early times, when stereoscopes where made widely available, because the production of stereo pairs become easier as photographic techniques evolved. \\cite{Darrah77} discusses these early ages (from 1851 to 1935) of stereoscopy in depth, and we refer the interested reader to his work for a detailed overview of the various applications of stereo pairs in those times.  Although the principle behind stereoscopy has remained the same ever since, there have been regular improvements to the methods of production and  visualization. In fact, the interest in stereoscopy has been closely linked to the development of both imaging and  visualization techniques \\citep{Okoshi76}, and peaks of interest arose as new production and/or  visualization tools were invented. Beside the evolution of photographic techniques, the advent of computers and their ability to produce accurate and detailed stereo pairs represents one such development which resulted in a peak of interest for stereoscopy that started in the 70'.  Several viewing devices have also been developed over the years, with one common aim: increasing the comfort and simplicity of stereoscopy for the viewer. In comparison to individual stereographs, the development of specialized glasses (red-blue, polarised, shutter-type) made stereo pairs easier to visualise. Lately, new stereoscopic technologies are being integrated into consumer products at a fast pace; movies, televisions, gaming consoles, cell phones, advertisement panels, and so on. Stereoscopy has become especially popular in the movie industry in the past few years with the advent of digital 3D cinemas. One should nonetheless not forget that stereo movies themselves are not recent: \\emph{The Power of Love}, in anaglyph 3D, first aired in 1922 \\citep[]{Zone07}.  In the scientific community, stereoscopy has been known and used in the past, but the extent to which it has been exploited in the presentation of data differs from field to field. In Astrophysics, stereoscopy has not been used extensively, despite the multi-dimensional nature of many data sets. Often, a data cube is sliced or projected in order to obtain 2D publishable pictures and graphs. The issue of displaying and publishing multi-dimensional data sets has been identified in the past, and some interesting (non-stereoscopic) solutions have been proposed. Cosmologists working on the time evolution of the large scale structures in the Universe using 3D movies to illustrate their simulations results is one example \\citep[e.g.][]{Holliman10}. Recently, \\cite{Barnes08} described how documents in an \\emph{Adobe Portable Document Format} (.pdf) are now able to contain animated 3D models, and described how this can be used to create interactive 3D graphs. In addition, \\cite{Barnes06} developed a 3D plotting library specifically tailored to the needs of the Astrophysics community. \\cite{Fluke06} also discuss and present alternative advanced image displays which might potentially take on a more significant place within the Astrophysics community in the future.  Yet, stereoscopic techniques are not unknown to Astrophysicists. Planetary scientists, for example, use red-blue anaglyphs. In the case of Mars \\citep[e.g.][]{Neukum04,Keszthelyi08}, several probes and remote sensing satellites were equipped with special stereo cameras, for example the European Space Agency \\emph{Mars Express} satellite and its High Resolution Stereo Camera Experiment \\citep[][]{Jaumann07}, the \\emph{Mars Reconnaissance Orbiter} and its High Resolution Imaging Science Experiment (HiRISE) camera \\citep[][]{McEwen07}, the \\emph{Phoenix} lander and its Surface Stereo Imager (SSI) \\citep[][]{Lemmon08} or the imager for the Mars Pathfinder (MPI) mission \\citep[][]{Smith97}.  Stereo pairs are another type of stereoscopic solution that has been employed to accommodate multi-dimensional data sets in publications. One of the first Astronomical stereo pairs published depicted the Moon and was created as early as 1862 by L.M. Rutherford \\citep{Darrah77}. By taking two subsequent images of the Moon with a six days interval, he obtained a strong enough change in orientation to induce a reasonable feeling of depth. More recently, the advent of computers expanded the possible applications of stereo pairs in Astrophysics. For example, \\cite{Yahil80} used them to display the position of the galaxies in the Revised Shapley-Ames Catalog; \\cite{Martinet81} and \\cite{Martinet81b} used stereo pairs to illustrate the shape of invariant surfaces in their study of dynamical problems with 3 degrees of freedom; \\cite{Weygaert89} and \\cite{Icke91} created stereo pairs to illustrate the Voronoi model they used to describe the asymptotic distribution of the cosmic mass on 10-200 Mpc scales; \\cite{Rhee91} produced stereo pairs to show a 3D map of their sample of Abell clusters; \\cite{Koutchmy92} published a stereo pair of the Solar Corona during the 1991 Solar eclipse; \\cite{Sofue94} illustrated the stripping of the LMC H{I} and molecular clouds; \\cite{Selman99,Selman99-3} created stereo pairs of the colour-magnitude diagrams of the ionizing cluster 30 Doradus; \\cite{Sirko04} used stereo pairs to display the 3D position of the blue horizontal-branch (BHB) stars discovered in the SDSS \\citep[][]{York00} spectroscopic survey; and \\cite{Vogt10a} and \\cite{Vogt10c} used stereo pairs in complement to their interactive 3D maps of the oxygen-rich material in SNR 1E 0102.2-7219 and SNR N132D. These examples do not represent an exhaustive list of all the work that has been published using stereo pairs in Astrophysics, but illustrate the wealth of topics that can profit and make use of this technique.  Nonetheless, the use of stereo pairs in Astrophysics is less prevalent than in other fields. In Biochemistry, for example, they have been an important tool to publish the 3D shapes of molecules from the beginning of the computer-era \\citep[e.g.][]{Hamilton70,Hardman73} until today \\citep[e.g][]{Pujadas01,Landsberg06, Xiang10}.  We believe that stereo pairs are a valuable tool in Astrophysics too, which recent 3D innovations may help renew. In other words, we argue that stereoscopy has a great but under-exploited potential for the publication of multi-dimensional Astrophysical data sets and can be a valuable complement to more standard plotting methods, especially with today's computing abilities. In Sec.~\\ref{Sec:real}, after introducing the \\emph{free-viewing} technique, we discuss the various theoretical ways to construct a stereo pair. In Sect.~\\ref{Sec:tools}, we present our trade-off method to efficiently create stereo pairs of data cubes using \\textsc{Python}. We then illustrate the unique features of stereo pairs as compared to standard plots in three different examples in Sect.~\\ref{Sec:science}: a conceptual one in Sect.~\\ref{Sec:conc}, one based on observational data in Sect.~\\ref{Sec:N132D} and one based on theoretical data in Sect.~\\ref{Sec:Alex}, for which the use of stereo pairs provides critical scientific information. We discuss the role that stereo pairs can play regarding future developments of 3D visualization techniques in Sec.~\\ref{Sec:future}, and summarize our conclusions in Sect.~\\ref{Sec:concl}. ",
        "conclusions": "\\label{Sec:concl}  We have discussed the concept of stereo pairs and highlighted their potential benefits for the Astrophysics community. First, we presented the free-viewing technique and provided advice to easily visualize both parallel and cross-eyed stereo pairs for the first time. We then argued that stereo pairs can be easily produced and reproduced on a computer screen or on printed material with most of the usual programming languages or software used nowadays within the Astrophysics community. In particular, we have introduced and described an alternative, off-the-shelf, easy way to produce high quality stereo pairs with \\textsc{Python}, which we refer to as the \\emph{simplified Toe-in} method. This technique adapts the \\emph{official} Toe-in procedure taking into account the current limitation of \\textsc{Python} plotting abilities, without affecting on the quality of the stereo pairs. Specifically, no vertical parallax can be detected in the resulting stereo pairs. Testing our sTi stereo pairs on several students and astronomers at the Mount Stromlo Observatory revealed that they represent a good trade-off, by being able to convey a satisfactory feeling of depth from any viewpoint, and by being as effective as standard Toe-in stereo pairs with an elevation lower that $|\\theta_0|\\sim50^{\\circ}$. The tests also revealed that sTi stereo pairs provide more depth structure around the data itself as compared to their equivalent Offset stereo pairs, and that the sTi method is in that respect more appropriate for creating stereo pairs in Astrophysics, a fact already observed by \\cite{Peterka09}.  We have then used three examples, one idealized and two realistic, with which we have presented various types of stereo pairs, highlighted several aspects of stereoscopic visualization, and identified the main benefits of using stereo pairs as a complement to more standard plotting techniques in a publication. First, they are a polyvalent tool that is adaptable to one's needs; their shape, size, and color can be adapted to best reveal the 3D data set without impacting the ability to transmit a depth perception to the viewer. Second, they profit any multivariate data set, observational or theoretical, and potentially benefit different genres of studies (e.g., of both the theoretical and observational kind). Third, they greatly facilitate the communication of complex 3D shapes. Especially, where a text description might be subject to interpretation, stereo pairs can force upon the viewer a unique view of the data set, thereby avoiding misconceptions. This is possibly the main factor that should dictate the use of stereoscopy in publications and presentations.  For all these reasons, stereo pairs should be considered a valuable tool for the Astrophysics community - a field where most data sets are multidimensional and multivariate, and where stereo pairs can be applied to many different sub-topics, but always with the common aim of simplifying, clarifying, and eliminating misconceptions. The evolution of informatics has made stereo pairs aesthetic, useful, and straightforward to produce. We are convinced that they have a promising future, given the rapid evolution of 3D visualization hardware and techniques, e.g. 3D televisions. Although we still lack a standardized API and user-friendly software to couple the stereo images to the display devices, these will likely be provided as soon as the need is identified. Sharing astrophysical data sets in 3D on hand-held devices might sound futuristic. Nonetheless, stereo pairs can already be easily stacked into a movie, and played during a talk in a lecture theatre equipped with 3D projection abilities, enabling the audience to experience a glimpse of the 3D future for Astrophysics. In conclusion, we are convinced that the ideas conceived through the ongoing 3D trend currently occurring in the non-scientific community can and ought to be used in Astrophysics. Stereo pairs are a good way to start opening our minds today."
    },
    "1107/1107.4419_arXiv.txt": {
        "abstract": "In order to reproduce the low mass end of the stellar mass function, most current models of galaxy evolution invoke very efficient supernova feedback. This solution seems to suffer from several shortcomings however, like predicting too little star formation in low mass galaxies at z=0. In this work, we explore modifications to the star formation (SF) law as an alternative solution to achieve a match to the stellar mass function. This is done by applying semi-analytic models based on De Lucia \\& Blaizot, but with varying SF laws, to the Millennium and Millennium-II simulations, within the formalism developed by Neistein \\& Weinmann. Our best model includes lower SF efficiencies than predicted by the Kennicutt-Schmidt law at low stellar masses, no sharp threshold of cold gas mass for SF, and a SF law that is independent of cosmic time. These simple modifications result in a model that is more successful than current standard models in reproducing various properties of galaxies less massive than $10^{10}{\\rm M}_{\\odot}$. The improvements include a good match to the observed auto-correlation function of galaxies, an evolution of the stellar mass function from $z=3$ to $z=0$ similar to observations, and a better agreement with observed specific star formation rates. However, our modifications also lead to a dramatic overprediction of the cold mass content of galaxies. This shows that finding a successful model may require fine-tuning of both star formation and supernovae feedback, as well as improvements on gas cooling, or perhaps the inclusion of a yet unknown process which efficiently heats or expels gas at high redshifts. ",
        "introduction": "\\label{sec:intro}  Ever since first introduced by  \\citet{white1991}, semi-analytic models of galaxy formation and evolution  \\citep[hereafter SAM;][]{kauffmann1999, kang2005, croton2006, Bower2006, cattaneo2006, DLB2007, monaco2007, somerville2008, Khochfar2009} have been successfully used to study how different physical processes determine the formation and evolution of galaxies. Based on halo merger trees extracted from $N$-body simulations or analytic methods, SAMs follow the main processes that are thought to affect the properties of galaxies, like gas cooling, star formation, feedback, and merging. These models provide a useful tool to study the interplay and the relative importance of these different physical processes. A detailed review of the semi-analytic method can be found in \\citet{baugh2006}.  Although various observational properties of galaxies are matched by the models, current SAMs still have problems in reproducing some important observations and up to now, there is no semi-analytic model that is able to fit all the key statistical properties of the observed galaxy population. For example, \\citet{guo2010}  show that the model of \\citet{DLB2007} substantially overproduces the low mass end of the stellar mass function of galaxies. This problem becomes more severe as the resolution of the underlying dark matter simulation increases. In addition, stellar mass functions at high redshifts are normally not well reproduced \\citep{fontanot2009, marchesini2008}; a good match to the 2-point auto-correlation function of galaxies is so far difficult to achieve  \\citep[e.g.][]{guo2010}; and the relation between the specific star formation rate and galaxy stellar mass deviates from observations  \\citep{somerville2008, fontanot2009}.  These discrepancies may have various reasons: inaccurate physical modeling of the processes that govern galaxy formation; technical problems in tuning the model against a large set of observational constraints; the use of a fixed functional form for a poorly understood process, which overly limits the freedom in tuning the model; or a wrong cosmological model adopted in the simulations. This large range of possibilities makes it difficult to correct identified discrepancies between model and observations.  For example,  \\citet{guo2010} tried to fix the over-prediction of the low mass end of the stellar mass function, as modeled by  \\citet{DLB2007}, by increasing the effect of supernovae feedback. They managed to reproduce the amplitude of galaxy stellar mass function, and also obtained a reasonable match to the galaxy luminosity functions in different bands. However, they also predict a too large fraction of red galaxies at low masses, too high amplitudes of the stellar mass functions in the redshift range of [0.8, 2.5], and galaxy auto-correlation functions that are too high for galaxies less massive than $6\\times10^{10}{\\rm M}_{\\odot}$. Moreover, the high feedback efficiency as used by  \\citet{guo2010} is physically difficult to motivate  \\citep{benson2003}, and is far more efficient than various solutions adopted by hydrodynamical simulations  \\citep{maclow1999, strickland2000, avilareese2011}.  In this work, we therefore explore an alternative solution. We tune the SF recipe instead of the supernovae (SN) feedback to study how our changes affect different statistics of galaxies, and to what degree the discrepancies mentioned above can be alleviated. Most of the current SAMs use an analogue to the empirical Kennicutt--Schmidt law \\citep{schmidt1959, kennicutt1998} to calculate the star formation rate (SFR) in galaxies. In this standard prescription, the SFR is roughly proportional to the cold gas mass, and scales inversely with the typical time-scale of a galactic disk \\citep[e.g.][]{cole2000}. This law is combined  with a sharp threshold at low gas densities, below which no SF occurs  \\citep{kauffmann1996, croton2006}. The use of a such a threshold is motivated both theoretically \\citep[e.g.][]{toomre1964,kennicutt1989, schaye2004} and observationally \\citep[e.g.][]{martin2001}.  This simple SF law is however likely an oversimplification. In recent years observational determinations of the SF law in galaxies have become increasingly refined, using various gas components (HI, CO) in combination with more reliable estimates of SF, based on UV and IR light. These recent findings can be summarized as follows: First, there are indications that the threshold for SF at low mass densities is not sharp. Instead, the SF efficiency drops off as a steep power-law at low gas densities \\citep{kennicutt2007, bigiel2008, wyder2009, roychowdhury2009,bigiel2010}. Second, several studies find that the SF rate is correlated more strongly with the mass of molecular gas (H$_2$) than with the the atomic gas \\citep{wong2002,bigiel2008,leroy2008}, probably also at high redshift \\citep{bouche2007,genzel2010}. This shows that simply correlating the star formation rate with the total cold gas density in models may not always lead to realistic results. It is also not clear which gas mass correlates best with the SF rate when averaging over the entire galaxy \\citep[e.g.][]{saintonge2011}. Third, there are both theoretical and observational indications that the normalization of the SF law might be lower at high redshift than locally \\citep[e.g.][]{wolfe2006, rafelski2010, gnedin2010, agertz2011, krumholz2011}. This means that simply extrapolating the local relation, as done in most SAMs, may be incorrect.  Several recent  models have attempted to address these issues. \\citet{baugh2005}, \\citet{weinmann2011a}, and \\citet{krumholz2011} specifically lower the quiescent star formation efficiencies at high redshifts in order to better match some observed properties of high redshift galaxies. Both  \\citet{fu2010} and  \\citet{lagos2010, lagos2011} focus on the first two of the above points, and present SAMs with updated and much more detailed SF recipes in comparison to previous SAMs. They do not include a sharp threshold for SF, and their SF rate depends on the molecular gas density instead of on the cold gas mass, following recent empirical and theoretical models \\citep[e.g.][]{blitz2006, krumholz2009a}. Their detailed models for SF depend on various internal properties of the disk, like size and pressure, which are not trivial to model in a SAM.  It is certainly worthwhile to try and include more detailed and observationally and theoretically better motivated SF law into SAMs. We present a complementary approach in this work, without taking into account complicated processes on sub-galactic scales, like the conversion from atomic to molecular gas. We instead try to solve in a straightforward way the inverse problem, namely which realistic SF law at galactic scales is required to improve the agreement between model galaxies and observations. We start with a standard model, similar to \\citet{DLB2007} and Neistein \\& Weinmann (2010), and assume that the SF rate depends on the cold gas mass in the galaxy, cosmic time, and in addition the host halo mass. We make simple change to the standard SF law that are qualitatively, but not quantitatively plausible, and explore how those impact on the properties of galaxies. We show that we can improve the agreement between SAMs and observations in several key aspects in this way.  We use the model developed by Neistein \\& Weinmann (2010), and implement it on both the Millennium Simulation  \\citep[MS,][]{springel2005} and the Millennium-II Simulation  \\citep[MS-II,][]{boylan2009}. The properties of galaxies with masses as low as $\\sim$ $10^{8}{\\rm M}_{\\odot}$ can probably be studied reliably with the help of these simulations \\citep{guo2010}, although this can be resolution dependent for extreme models (see Neistein \\& Weinmann 2010). Our aim is to investigate how much a model based on \\citet{DLB2007} can be changed and improved by tuning the SF law alone. We thus do not change gas cooling and feedback in our models, but leave it at the default standard values (except in one case, for illustration, as will be explained).    This paper is organized as follows. In section \\ref{sec:model} we present the different models used in this work. Our starting point is a model that is based on the widely used SAM of \\citet{DLB2007}, with an improved prescription for hot gas stripping of satellite galaxies. We then develop four different models. Of these, models 2 and 3 include the key changes to the quiescent and burst mode star formation. In section \\ref{sec:results} we present predictions for the stellar mass functions at low and high redshifts, the relation between galaxy specific SF rate and stellar mass, the galaxy cold gas mass function, the auto-correlation functions, and the SF rate density as a function of redshift. A discussion of the implications of our results, and the conclusions are presented in section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We use the method developed by  \\citet{NW2010}, combined with both the Millennium (MS) and Millennium-II (MS-II) cosmological simulations, to study the effect of modifying the star formation (SF) recipe in low mass galaxies. We show that by modifying SF in both the quiescent and the burst mode, the stellar mass function observed in the local Universe can be reproduced well down to $10^{8.5}{\\rm M}_{\\odot}$. Simultaneously, the models can fit the observed median SSFR--$M_{\\rm star}$ relation for galaxies less massive than $10^{10}{\\rm M}_{\\odot}$, the correlation functions for galaxies more massive than $10^{8.77}{\\rm M}_{\\odot}$, the stellar mass functions up to redshift of around $3$, and the general trend of the SFR density as a function of redshift.  The modifications to the SF recipe in our models with respect to standard SAMs (e.g. DLB07) include: \\begin{enumerate} \\item no sharp threshold in the cold gas mass for SF; \\item letting the SF efficiency in the quiescent mode depend on host halo mass; \\item removing the dependence of the star formation rate on Hubble time (which came via the dependence on disk dynamical time); \\item a lower star burst efficiency in low mass haloes; \\item additional modifications of cooling and SF in massive haloes to match the properties of high mass galaxies. \\end{enumerate}  Model 2) only includes the changes to the quiescent mode of SF, i) - iii). We show that this is not enough to reproduce the low mass end of the SMF, as star formation in the burst mode compensates for the changes to the SF law in the quiescent mode. This is also the reason why \\citet{lagos2010} have found that their changes to the SF law in the quiescent mode does not change the resulting SMF much. In model 3), we have thus additionally decreased SF in the burst model (modification iv), which results in a clearly improved SMF. We also considered a model 3b) which is similar to model 3), except that we allow for the usual time-dependence of the SF law (i.e. do not make modification iii). Results of model 3b) are similar to model 3) out to $z \\sim 2$, except that the SSFR--$M_{\\rm star}$ relation is slightly less well reproduced. It is thus not clear whether modification iii) is necessary.  Note that removing the time-dependence of the star formation efficiency is a significant change with respect to previous models. Even in the recent models of \\citet{fu2010} and \\citet{lagos2010,lagos2011}, where SF efficiency does not depend on the disk dynamical time of galaxies, a time-dependence enters via the conversion efficiency from atomic to molecular gas, that depends on the gas density. In order to justify the behaviour we suggest in model 3), we would need to postulate a mechanism that scales with time in the opposite way than usually assumed, like for example a metallicity-dependent conversion of atomic to molecular gas \\citep{krumholz2011}. We note that a weak dependence of the SF efficiency on cosmic time was seen in hydrodynamical simulations \\citep{neistein2011c}, although the reason for this behaviour is still not clear.  In our models 2) -- 4), the ratio between the quiescent star formation rate and the cold gas mass, $\\dot{M}_{\\rm star}/M_{\\rm cold}$ which is equal to the gas consumption timescale, is roughly proportional to $M_{\\rm halo}^{2}$ for low mass haloes and is almost independent of halo mass for high masses. This may in fact be supported by recent observations. Recently, \\citet{shi2011} derive an extended Schimdt law from an observed galaxy sample that extends over 5 orders of magnitude in stellar density, including galaxies with low surface brightness. They find that $\\dot{M}_{\\rm star}/M_{\\rm cold}$ is proportional to $M_{\\rm star}^{0.52}$, with a 1-$\\sigma$ scatter of $0.4$ dex. The stellar mass of galaxies has been claimed to obey a tight relation with the host halo mass \\citep{conroy2009, moster2009, guo2010a}. It is proportional to $M_{\\rm halo}^{2.8}$ for low mass haloes and to $M_{\\rm halo}^{0.2}$ at high mass end \\citep{ wang2006}. Without considering the scatter of the relation, this indicates that the observed $\\dot{M}_{\\rm star}/M_{\\rm cold}$ is roughly proportional to $M_{\\rm halo}^{1.46}$ at low masses and $M_{\\rm halo}^{0.1}$ at high mass end. The dependence on halo mass in our models is therefore quite close to the observational result by  \\citet{shi2011}.  In all our models, the cold gas mass function of galaxies is dramatically over-predicted. This is because in DLB07, the total amount of cold gas and stellar mass exceeds the observed total amount \\citep[see][]{obreschkowCroton2009, fu2010}. Decreasing SF rates in low mass galaxies while letting cooling and feedback recipes remain unchanged naturally results in an overproduction of the cold gas mass in low mass galaxies. This is not only a problem for the models we present here. As shown recently by \\citet{lu2011}, when the model K-band luminosity function is forced to fit the data in the local Universe, the cold gas mass function is dramatically over-predicted in all the semi-analytic models they study. This is consistent with our models over-predicting the cold gas mass functions when we fit the stellar mass function at z=0. On the other hand, with similar cooling and supernovae feedback recipes, the models that do fit the cold gas mass function in turn have problems in reproducing the stellar mass functions \\citep[e.g.][]{fu2010,obreschkowCroton2009}. In the meantime, approaches like increasing the SN feedback, that suppress the total amount of cold gas and stellar mass, lead to several other serious problems \\citep{guo2010}. These results show again that it is currently difficult for a single model to fit all observations, as mentioned in Sec.~\\ref{sec:intro}, unless we allow for free tuning of all recipes, including cooling (see NW2010).  Although the modifications of the SF law presented in this work help to improve the agreement with several observed statistical properties of galaxies, and seem to follow a similar scaling like the observed SF efficiencies in galaxies, the normalization of the SF efficiency cannot be correct. This is clear from the fact that our models overpredict the cold gas mass function. With a lower cold gas fraction in galaxies, the SF efficiency would obviously have to be higher than currently assumed in the models, to obtain the same SF rate.  In Fig.~\\ref{fig:SFlaw}, we compare the SFR--HI mass relation in model 3) combined with MS-II (right panel), with a recent observation of the FUV derived SFR--HI mass relation for galaxies within $\\sim 11$ Mpc of the Milky Way \\citep[][left panel]{lee2011}. The cold gas mass in model 3) is converted to the HI mass to be compared with observation, using a correction including three factors. First, as for Fig.~\\ref{fig:CMF}, the cold gas mass from model 3) is divided by a factor of $1.45$, to account for a warm ionized gas phase, as done in \\citet{obreschkowCroton2009}. Second, we multiply the results by a factor of $0.76$  to remove the contribution of helium and heavier elements  \\citep{power2010}. Finally, the hydrogen gas mass is divided by $1.4$, as adopted by \\citet{power2010} and \\citet{lu2011}, to remove the contribution of $H_2$. Since the galaxies in the observed sample have stellar masses less than $\\sim 10^{10}{\\rm M}_{\\odot}$, we present model results for galaxies with $8<\\log (M_{\\rm star})<10$. The median relation from the observations is plotted as red line in each panel. Although the observations of \\citet{lee2011} are limited to a small volume of space, the obviously different relations in observations and in model 3) indicate that the SF efficiency in model 3) is indeed much lower than in reality.  \\begin{figure*} \\bc \\hspace{-1.4cm} \\resizebox{14cm}{!}{\\includegraphics{./SFlaw.ps}}\\\\% \\caption{SF rate -- HI gas relation in model 3) compared with observation. The left panel is the FUV SFR--HI mass results of \\citet{lee2011}. The right panel shows results of 3) combined with MS-II simulations, for galaxies with $8<\\log(M_{\\rm star})<10$. Gray scales correspond to the logarithm of the number of galaxies. Red line in each panel is the median relation of the observed galaxies of  \\citet{lee2011}. } \\label{fig:SFlaw} \\ec \\end{figure*}   Fitting the most important observed properties of galaxies is not a trivial task. It is not clear up to now if this difficulty in modeling galaxy properties reflects a fundamental problem in our understanding of the dark matter universe, or if it is mainly due to an insufficient understanding of the baryonic physics involved in galaxy formation and evolution. For example, perhaps SAMs miss an important ingredient of galaxy formation, like a mechanism that preheats the gas in the universe so that it cannot cool to low mass haloes \\citep[][but see also \\citet{crain2007}]{mo2005}, or a form of feedback that mainly heats low entropy gas in high redshift haloes \\citep{McCarthy2011}. Alternatively, it could be that cooling is over-efficient in the current SAMs for some reason.  With the tests carried out in this work, we show that only modifying star formation in the SAMs can already improve agreement with observations in several important aspects. Up to now, a high-resolution SAM that matches the SSFR-stellar mass relation at $z=0$, the SMF and its evolution, the correlation functions and the cold gas fractions of galaxies simultaneously, does not yet seem to exist. To find such a model, approaches that allow scanning of a large parameter space \\citep{henriques2009,lu2010,bower2010}, and approaches that allow deviations from the usually assumed functional forms for physical recipes \\citep{NW2010}, or that include other physical processes than currently considered \\citep{henriques2010} may be promising. However, even if one or several models are found that do indeed reproduce all these fundamental observables, it will be important to identify degeneracies, and to verify whether the models are physically plausible and can be brought into agreement with alternative approaches, like predictions from hydrodynamical simulations."
    },
    "1107/1107.4133_arXiv.txt": {
        "abstract": " ",
        "introduction": "}  The classification of star clusters as either ``open'' or ``globular'' may first have been explicitly adopted by \\citet{Shapley1916}. However, it is clear that neither the ``classical'' globular nor open clusters constitute homogeneous classes of objects, and furthermore the boundary between the two classes is somewhat fuzzy. The Galactic globular clusters (GCs) can be grouped into at least two sub-populations, loosely associated with the Galactic halo and bulge/thick disc, that have distinct kinematics, spatial distributions and metallicities \\citep{Zinn1985,Minniti1996}. Even in our Galaxy, the open clusters span large ranges in age, mass and concentration that overlap somewhat with the globular clusters, and making a clear-cut distinction between open and globular clusters becomes increasingly difficult as one moves to external galaxies \\citep[e.g.][]{Hodge1988}. This has become particularly true in the last few decades due to the discovery of young star clusters with masses in the range $10^5 - 10^7 M_\\odot$, predominantly in interacting galaxies and starbursts \\citep{Whitmore2003,PortegiesZwart2010}, but also in some normal non-interacting spirals \\citep{Larsen2000a}. To accommodate these objects, terms such as young massive clusters (YMCs), super-star clusters (SSCs), etc., have been introduced.  Although classification is often a natural first step in the study of a new field, there is no generally accepted classification scheme for star clusters that has managed to gain foothold, and it is not clear to what extent the different labels that have been adopted by various authors correspond to physically different objects \\citep[see an excellent discussion of this point in][]{Terlevich2004}. It therefore appears more useful to discuss the differences and similarities among cluster systems in different environments in terms of quantifiable parameters. To this end, it is useful to distinguish between three different aspects of a cluster system: 1) The formation efficiency of (long-lived) star clusters relative to field stars, 2) the shape of the initial cluster mass function (ICMF), and 3) dynamical effects. Items (1) and (3) are dealt with extensively elsewhere in these proceedings (e.g.\\ Bastian, Bressert, Gieles, Kruijssen). Here I will mainly concentrate on (2), the ICMF. ",
        "conclusions": "Comparison with other nearby spiral galaxies suggests that the Milky Way should contain a significant population of young star clusters with masses above $10^5 M_\\odot$. Scaling from the cluster systems of NGC~5236 and NGC~6946, about 20 clusters with $M>10^5 M_\\odot$ and ages younger than 200 Myr are expected. Are these objects young analogues of the ancient globular clusters that populate the halo and bulge of our own and other galaxies? In terms of masses and sizes, the young massive clusters in nearby spirals, as well as in on-going starbursts and mergers, certainly appear similar to their older counterparts. Differences in the ``typical'' masses may be largely due to a combination of selection effects and dynamical evolution -- the ``initial cluster mass function'' (ICMF) appears to extend well above $10^5 M_\\odot$ whenever there are enough clusters to sample the high-mass end, although there appears to be an environmentally dependent truncation that occurs at a few times $10^5 M_\\odot$ in present-day, quiescent discs, but at $>10^6 M_\\odot$ in starbursts and ancient GC populations.  A more significant difference between ancient GCs and present-day cluster formation may be the presence of multiple stellar populations in the former. These are now revealed by direct colour-magnitude diagrams, although the presence of abundance anomalies have hinted at this for many decades. It remains unclear, however, whether these phenomena are uniquely related to ancient GCs. The surviving GCs must have had initial masses well above $10^5 M_\\odot$, and little is currently known about the detailed star formation histories of young clusters in this mass range, simply because they are generally too distant to observe individual stars in detail. However, there is evidence from CMDs that some young clusters in the Large Magellanic Cloud, as well as in more distant galaxies, may have extended star formation histories, although this is not entirely conclusive yet.  \\small  %"
    },
    "1107/1107.1716_arXiv.txt": {
        "abstract": "Baryon-density perturbations of large amplitude may exist if they are compensated by dark-matter perturbations such that the total density is unchanged.  Primordial abundances and galaxy clusters allow these compensated isocurvature perturbations (CIPs) to have amplitudes as large as $\\sim10\\%$. CIPs will modulate the power spectrum of cosmic microwave background (CMB) fluctuations---those due to the usual adiabatic perturbations---as a function of position on the sky.  This leads to correlations between different spherical-harmonic coefficients of the temperature and/or polarization maps, and induces polarization $B$-modes.  Here, the magnitude of these effects is calculated and techniques to measure them are introduced. While a CIP of this amplitude can be probed on large scales with existing data, forthcoming CMB experiments should improve the sensitivity to CIPs by at least an order of magnitude. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.3978_arXiv.txt": {
        "abstract": "Small fraction of isocurvature perturbations may exist and correlate with adiabatic perturbations in the primordial perturbations. Naively switching off isocurvature perturbations may lead to biased results.  We study the effect of dark matter isocurvature on the structure formation through N-body simulations.  From the best fit values, we run four sets of simulation with different initial conditions and different  box sizes. We find that, if the fraction of dark matter isocurvature is small, we can not detect its signal through matter  power spectrum and two point correlation function with large scale  survey. However, the halo mass function can give an obvious signal. Compared to $5\\%$ difference on matter power spectrum, it can get $37\\%$ at $z = 3$ on halo mass function. This indicates that future high precise cluster count experiment can give stringent constraints on dark matter isocurvature perturbations. ",
        "introduction": "Recent high accuracy observations, such as  cosmic microwave background radiation (CMB)\\cite{wmap7}, and the large scale structure (LSS)\\cite{sdss}, provide us wealthy information on the universe. Currently, so called concordance cosmological model, in which approximately scale-invariant, Gaussian, adiabatic  primordial perturbations seed the structure of our universe, is mostly favored.  However, there exist  some models which predict non-negligible  isocurvature perturbations,  such as multiple scalar fields inflation\\cite{Liddle:1998jc,  Kanti:1999vt, Copeland:1999cs, Langlois:1999dw, Wands:2002bn, Piao:2002vf, Dimopoulos:2005ac, Langlois:2008wt, Cai:2009hw, Cai:2008if, Cai:2010wt} and curvaton\\cite{Lyth:2001nq,Moroi:2001ct,  Lyth:2002my} models. The primordial isocurvature perturbations then can induce the cosmological matter  perturbations such as cold dark matter (CDM), baryons, dark energy and neutrino\\cite{Seckel:1985tj,  Linde:1991km, Linde:1996gt,Liu:2010ba} in the radiation era.  Although the pure isocurvature perturbations have already been ruled out by the observation of Boomerang and MAXIMA-1\\cite{Enqvist:2000hp}, models  with primordial perturbations comprised of dominate adiabatic and a small fraction of isocurvature modes\\cite{Gordon:2002gv, Crotty:2003rz,  Bucher:2004an, Moodley:2004nz, KurkiSuonio:2004mn, Beltran:2005xd, Bean:2006qz, Trotta:2006ww, Sollom:2009vd,  Valiviita:2009bp, Beltran:2005gr,  Mangilli:2010ut, Li:2010yb} still survive. Due to the quality of data, we still  can not confirm  the  destiny of the isocurvature perturbation.  Moreover, if we  switch off the isocurvature perturbations naively, the result would be misleading\\cite{Zunckel:2010mm, Liu:2010ba}. So, it is important to seek a way  to detect the isocurvature perturbations and give stringent constraints on isocurvature perturbations parameters.  In  Ref. \\cite{Li:2010yb}, we have given constraints on parameters of two different cosmological models  with latest astronomical data.  One is with pure adiabatic initial condition (IC) and the other is with mixed IC in which  small fraction of isocurvature mode is correlated with dominating adiabatic mode through $\\cos\\Delta$. We found that compared to adiabatic mode, the isocurvature perturbations have smaller amplitude ($A_s^{iso}/A_s^{adi} \\sim 10^{-2}$) and bluer tilt ($n_s^{iso} \\sim 2$). That is to say the isocurvature modes is pronounced on small scales ($k > 1\\kunit$). However, on these scales, the structure grows non-linearly, so we can not study the details of evolution in a full analytical way and then test them against future high-precision observations.  On the other way, N-body simulations plays more and more important role in cosmology nowadays, especially in the study of nonlinearity. It can make predictions and compare with observations for a specific model. It can also be used to check the validity of a particular  method \\cite{Bagla:2004au}.  In this work, we study the effects of CDM isocurvature perturbations on LSS  by implementing N-body simulations. Given the best fit parameters in Ref. \\cite{Li:2010yb}, we carry out four sets of simulations with different ICs in different boxes. We find that,  unlike the power spectrum and two point correlation function, mass function is sensitive to dark matter isocurvature perturbations. This indicates  that, we can give stringent constraints on dark matter isocurvature perturbations  with future high  precision cluster counts experiment.  The structure of this paper is organized as follows. In Sec. \\ref{sec:ic}, we  briefly introduce the CDM isocurvature perturbation and set the ICs for N-body simulation. In Sec. \\ref{sec:nbody},  we describe the details of our N-body simulation and present the result. We give summary in Sec. \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} Isocurvature perturbations, inevitably generated from multi-field  inflation  or curvaton models, can be used to test these models.  Although pure isocurvature perturbation models have already been ruled out,  there still exists possibility that a  small fraction of isocurvature mode is correlated to  the dominating adiabatic perturbation.  This is important to the parameter estimation, since rough ignorance would lead to biased result.  With the best fit values obtained in Ref.\\cite{Li:2010yb},   we perform four sets of N-body simulations to seek a best way to detect the isocurvature perturbations.  We find that, if the fraction is small enough($A_{iso}/A_{adia}\\sim 3\\%$)  we can not peek it in the BAO  observation. The position and the width of the bump in two-point  correlation function,  which  mainly depend on the background parameter, are almost the same.  There are some differences in matter power spectrum and halo mass function. However, the  $5\\%$ difference in matter power  spectra  makes it hard to be observed.  On the contrary,  the deviation in the halo mass function is obvious. The difference is getting larger as we go to  higher redshift. For $M \\sim 5\\times10^{10} M_\\odot$ the discrepancy can get $37\\%$ at $z = 3$.   This implies that, with future precise cluster number count observations,  we can detect the initial condition of dark matter isocurvature and give stringent constraints."
    },
    "1107/1107.4844_arXiv.txt": {
        "abstract": "The monitor of all-sky X-ray image (MAXI) Gas Slit Camera (GSC) on the International Space Station (ISS) detected a gamma-ray burst (GRB) on 2009, September 26, GRB\\,090926B. This GRB had extremely hard spectra in the X-ray energy range. Joint spectral fitting with the Gamma-ray Burst Monitor on the Fermi Gamma-ray Space Telescope shows that this burst has peculiarly narrow spectral energy distribution and is represented by Comptonized blackbody model.  This spectrum can be interpreted as photospheric emission from the low baryon-load GRB fireball.  Calculating the parameter of fireball, we found the size of the base of the flow $r_0 = (4.3 \\pm 0.9) \\times 10^{9} \\, Y^{\\prime \\, -3/2}$ cm and Lorentz factor of the plasma $\\Gamma = (110 \\pm 10) \\, Y^{\\prime \\, 1/4}$, where $Y^{\\prime}$ is a ratio between the total fireball energy and the energy in the blackbody component of the gamma-ray emission. This $r_0$ is factor of a few larger, and the Lorentz factor of 110 is smaller by also factor of a few than other bursts that have blackbody components in the spectra. ",
        "introduction": "A discovery of the afterglows of gamma-ray bursts (GRBs) made it clear that GRBs are at cosmological distance and emit enormous energy.  The spectra of GRB prompt emission have been expressed empirically with a smoothly broken power-law model (``Band Function'') \\citep{1993ApJ...413..281B}. In order to explain the huge energy and non-thermal spectral shape, models of synchrotron emission from shock-accelerated electrons in the relativistic outflow \\citep[e.g.][]{1994ApJ...430L..93R} are suggested. However, sometimes the low-energy component is very hard and the power-law photon index $\\alpha$ \\footnote{We use the photon index $\\alpha$ in the context of $E^{\\alpha}$, where $E$ is the energy, throughout this paper.} becomes greater than the theoretical limit of $2/3$ \\citep{2002ApJ...581.1248P}. This observational fact has been explained in many ways such as non-thermal \\citep[e.g.][and references therein]{2009ApJ...702L..91M, 2010ApJ...725.1137L} and thermal \\citep[e.g.][]{2000ApJ...530..292M, 2004ApJ...614..827R, 2010MNRAS.407.1033B} models, but reasons for such hard spectra have not been understood clearly. For example, \\citet{2003A&A...406..879G} suggest that some thermal models are consistent with observed spectral characteristics of several GRBs with extremely hard spectra. On the other hand, \\citet{2005PASJ...57.1031S} remark that ``jitter'' radiation \\citep[e.g.][]{1973ApJ...183..611E, 2000ApJ...540..704M} is one of possible mechanisms of reproducing spectral indices $\\alpha>-2/3$.  A GRB was triggered by Swift/BAT and Fermi/GBM at 21:55 on 2009 September 26. The quick results of this burst, GRB\\,090926B, are reported to GCN by the Swift team \\citep{2009GCNR..246....1G} and the Fermi team \\citep{2009GCN..9957....1B}. Both of them conclude that the burst has a hard spectrum, whose photon index is $\\alpha > -2/3$, above the critical value of synchrotron shock models, so-called ``line of death''.  GRB\\,090926B was also observed with MAXI/GSC. MAXI/GSC can examine a low energy portion of the spectra of GRBs below 10 keV, while Swift/BAT or Fermi/GBM observes $\\gtrsim$ 10 keV.  The observational results below 10 keV may give even more severe constraints to interpretation of the spectra. In this paper, we report the observational results obtained with MAXI/GSC on GRB\\,090926B, and discuss their interpretation. ",
        "conclusions": "MAXI GSC observed the first 30 s of GRB\\,090926B prompt emission. From the data of the scans before and after the burst, we could not find any signal of emissions with the flux limit of about 3 $\\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$ (4--10 keV) for each scan. The joint spectral analysis with Fermi GBM reveals that the spectrum of GRB\\,090926B shows a peculiar narrow shape. The spectral index $\\alpha$ of time averaged spectrum is positive. The $E_{\\mathrm peak}$ of the burst is low relative to other bursts with such a hard spectral indices. This hard spectral index cannot be realized by interstellar absorption, synchrotron self absorption, synchrotron self Compton, nor jitter radiation. We find that the spectrum can be fit well by Comptonized blackbody model. The blackbody radiation can be interpreted as a photospheric emission of the GRB fireball.  Following the model by \\citet{2007ApJ...664L...1P}, we obtain the size of the base of the flow $r_0 = (4.3 \\pm 0.9) \\times 10^{9} \\,  Y^{\\prime \\, -3/2}$ cm and Lorentz factor of the plasma $\\Gamma = (110 \\pm 10) \\, Y^{\\prime \\, 1/4}$. According to \\citet{2010arXiv1011.6005B}, the observed photospheric spectrum is blackbody if the outflow energy is dominated by radiation rather than baryon up to the photospheric radius. Therefore, the spectrum of GRB\\,090926B may be the example of the low baryon-load fireball.  \\bigskip  This research was partially supported by the Ministry of Education, Culture, Sports, Science and Technology (MEXT), Grant-in-Aid No.19047001, 20041008, 20244015, 20540230, 20540237, 21340043, 21740140, 22740120, and Global-COE from MEXT ``The Next Generation of Physics, Spun from Universality and Emergence'' and ``Nanoscience and Quantum Physics''."
    },
    "1107/1107.3180_arXiv.txt": {
        "abstract": "We present results from the Radio Interferometric Planet (RIPL) search for companions to the nearby star GJ 896A. We present 11 observations over 4.9 years.  Fitting astrometric parameters to the data reveals a residual with peak-to-peak amplitude of $\\sim 3$ mas in right ascension.  This residual is well-fit by an acceleration term of $0.458 \\pm 0.032$ mas/y$^2$.  The parallax is fit to an accuracy of 0.2 mas and the proper motion terms are fit to accuracies of 0.01 mas/y.  After fitting astrometric and acceleration terms residuals are 0.26 mas in each coordinate, demonstrating that stellar jitter does not limit the ability to carry out radio astrometric planet detection and characterization. The acceleration term originates in part from the companion GJ 896B but the amplitude of the acceleration in declination is not accurately predicted by the orbital model. The acceleration sets a mass upper limit of 0.15 $M_J$ at a semi-major axis of 2 AU for a planetary companion to GJ 896A.  For semi-major axes between 0.3 and 2 AU upper limits are determined by the maximum angular separation; the upper limits scale from the minimum value in proportion to the inverse of the radius.  Upper limits at larger radii are set by the acceleration and scale as the radius squared. An improved solution for the stellar binary system could improve the exoplanet mass sensitivity by an order of magnitude. ",
        "introduction": "The study of extrasolar planets provides an important link between the study of star formation, the study of our own solar system, and the search for extraterrestrial life. \u00a0Radial velocity searches \\citep{2006ApJ...646..505B}, transit searches with Kepler \\citep{2011ApJ...728..117B}, and microlensing searches \\citep{2008Sci...319..927G} have surveyed significant components of the exoplanet mass-separation parameter space exposing a wide variety of phenomenology.  Despite their abundance, low mass stars have been less well-characterized for a number of reasons. In particular, low mass stars are often too faint or too strongly variable for high precision radial velocity and transit searches. Microlensing searches have identified companions to low mass stars \\citep{2008ApJ...684..663B,2006Natur.439..437B,2006ApJ...644L..37G}, which suggest that such companions may be very common.  On the other hand, radial velocity searches suggest very low planet occurrence rates for close-in companions \\citep{2006ApJ...649..436E,2007ApJ...670..833J}.  The frequency and distribution of companions to low mass stars can be an important diagnostic of methods for planet formation \\citep{2004ApJ...612L..73L,2005ApJ...626.1045I,2006ApJ...643..501B,2007Ap&SS.311....9K}.  Astrometry can be a powerful tool for detection and characterization of nearby, low-mass planetary systems. \\citet{2010AJ....140.1657M} demonstrate that optical astrometric measurements with the Palomar Testbed Interferometer are capable of achieving sub-milliarcsecond accuracy in the separation between close binaries.  They summarize the known properties of sub-stellar companions to low mass stars in binary systems and conclude that giant planet companions in such systems are common. The Carnegie Astrometric Planet Search Program \\citep{2009PASP..121.1218B} has also obtained sub-milliarcsecond astrometry on the low mass star NLTT 48256 over 2 years of observations; this program will ultimately survey 100 nearby M, L, and T dwarfs.  The Very Long Baseline Array (VLBA) can routinely achieve an astrometric accuracy of $\\sim100$ $\\mu$as in a single epoch at frequencies $\\lsim 10$ GHz and $\\sim 20$ $\\mu$as at higher frequencies. The VLBA is capable of accuracies as high as 8 $\\mu$as under favorable circumstances \\citep{2003ApJ...598..704F}. \u00a0 To use this technique the target source must have a sufficiently high brightness temperature to be detected by a high resolution radio interferometer. \u00a0For the VLBA, brightness temperatures must be $T_b \\gsim 10^{7}$ K, requiring nonthermal emission. \u00a0Thus, the systems which can be studied are limited to the most active stars. VLBA astrometry has been important for a wide range of stellar measurements including parallax \\citep[e.g.,][]{2007ApJ...667.1161S,2007A&A...474..515M,2007ApJ...671..546L,2009ApJ...698..242T}, galactic structure  \\citep{2006Sci...311...54X,2011ApJ...733...25X,2011AN....332..461B}, binary orbit characterization \\citep{2006A&A...446..733G}, and the radio/optical reference-frame link \\citep{1999A&A...344.1014L}. Recently, \\citet{2009ApJ...706L.205F} detected a brown dwarf with the VLBA and explored the possibilities of companion detection. \\citet{2007arXiv0704.0238B} reviewed the prospects for future radio astrometric surveys with respect to other methods.  RIPL, the Radio Interferometric Planet Search, is a 4-year program using the VLBA and the 100m Green Bank Telescope (GBT) to search for the astrometic signatures of massive planets around nearby ($D < 10$ pc), low-mass (M dwarf) stars. We are observing a sample of 30 stars a total of 12 times over a 4-year period with the expectation of achieving $\\sim 0.1$ milliarcsecond astrometric accuracy per epoch, and ultimately having the sensitivity to detect Jupiter-mass planets at 1 AU.  This astrometric search is an important complement to radial velocity, transit, microlensing, and optical astrometric searches for planets on account of greater sensitivity to  1) low mass stars, 2) long-period planets, 3) active stars, and, 4) planets that may be imaged with extreme adaptive optics. Many RIPL stars are in binary systems, which may have higher planetary fractions.  In Paper I \\citep{2009ApJ...701.1922B}, we used Very Large Array (VLA) observations to identify the sample of stars appropriate for RIPL and we used the VLBA to characterize the detectability and short-term ($\\lsim 10$ days) stability of the astrometric images of these stars.  The short-term stability addresses a significant concern of radio astrometry that stellar activity will lead to large astrometric errors.  The absence of a significant difference in the historical optical proper motions and the radio proper motion from this RIPL precursor survey permitted us to exclude brown dwarf companions to the four stars observed.  RIPL observing began in October 2007 and is anticipated to end in late 2011.  In this paper, we present initial results from RIPL for the star GJ 896A. This star was also observed as part of the precursor survey in paper I. Together these observations demonstrate astrometric stability over a 4.9 year span. GJ 896A (EQ Peg A) is a variable star in a multiple system with GJ 896 B, C, and D. GJ 896B is its closest companion; the orbit has been characterized, but it is not well constrained \\citep{1984AJ.....89.1063H}. GJ 896 A and B have spectral types M3.5 and M4.5.  These two stars have been detected previously in the radio and seen to have point-like structure and strong radio variability \\citep{1985A&A...149...95P,1989A&A...210..284J,1995A&A...298..187B,1998ASPC..154.1484G}. In \\S 2, we describe our observations, analysis, and results.  In \\S 3, we present our astrometric analysis of the data.  In \\S 4, we examine the constraints that we can place on companions to GJ 896A.  In \\S 5, we summarize our results. ",
        "conclusions": "We have presented here an analysis of RIPL observations of the star GJ 896A.  We measure the parallax to an accuracy of 0.2 mas and the proper motion to an accuracy of 0.01 mas/y.  We measure an acceleration term of $0.466 \\pm 0.045$ mas/y$^2$.  This term is significantly less than expected based on the model for the orbit of the GJ 896AB system, suggesting an error in the optically-determined orbital parameters.  Attributing this acceleration to a planetary mass companion, we demonstrate sensitivity to a planetary companion at 2 AU of mass 0.15 $M_J$.  The residuals that we obtain from our fitting of 0.25 mas in each coordinate over 4.9 y indicate the feasibility of planet searching and characterization with radio astrometry of nearby low mass stars.  These residuals can potentially be improved for this star and others through technical improvements to the data analysis including accounting for the motion of the star during the observing epoch and more exact removal of tropospheric fluctuations.  The residuals indicate that astrometric jitter from stellar activity on scales of a stellar radius or larger is not a limiting factor in planet astrometry.  Improvements in our analysis of systematic errors from stellar motion and tropospheric phase errors will permit us to explore sub-stellar-radius astrometric accuracy.  Future RIPL papers will explore analysis of additional data from this star, improved analysis techniques, and analysis of other stars in the sample.  Upgrades to the recording bandwidth of the VLBA (and GBT) will produce significant improvements in the capability to carry out observing programs of this kind.  The improved sensitivity will allow more rapid detection of bright stars such as GJ 896A, thereby reducing any smearing effect from stellar motion, allow the use of closer calibrators for improved tropospheric corrections and reduced time on calibrator, provide higher signal to noise ratio detections of fainter stars, and enable us to detect a larger sample of stars."
    },
    "1107/1107.4243_arXiv.txt": {
        "abstract": "{Intermediate-mass black holes (IMBHs) are of interest in a wide range of astrophysical fields. In particular, the possibility of finding them at the centers of globular clusters has recently drawn attention. IMBHs became detectable since the quality of observational data sets, particularly those obtained with HST and with high resolution ground based spectrographs, advanced to the point where it is possible to measure velocity dispersions at a spatial resolution comparable to the size of the gravitational sphere of influence for plausible IMBH masses.} {We present results from ground based VLT/FLAMES spectroscopy in combination with HST data for the globular cluster NGC 6388. The aim of this work is to probe whether this massive cluster hosts an intermediate-mass black hole at its center and to compare the results with the expected value predicted by the $M_{\\bullet} - \\sigma$ scaling relation.} {The spectroscopic data, containing integral field unit measurements, provide kinematic signatures in the center of the cluster while the photometric data give information of the stellar density. Together, these data sets are compared to dynamical models and present evidence of an additional compact dark mass at the center: a black hole.} {Using analytical Jeans models in combination with various Monte Carlo simulations to estimate the errors, we derive (with $68 \\% $ confidence limits) a best fit black-hole mass of $ (17 \\pm 9) \\times 10^3 M_{\\odot}$ and a global mass-to-light ratio of $M/L_V = (1.6 \\pm 0.3) \\ M_{\\odot}/L_{\\odot}$.} {} ",
        "introduction": "For a long time, only two mass ranges of black holes were known. On the one hand, we have stellar mass black holes, which are remnants of massive stars, and can be observed in binary systems. On the other hand, there are supermassive black holes at the centers of galaxies, some of them accreting at their Eddington limit and producing the brightest objects known (quasars).  It has been demonstrated that supermassive black holes show a tight correlation between their mass and the velocity dispersion of the galaxy in which they reside \\citep[e.g.][]{ferrarese_2000,gebhardt_2000, gultekin_2009}. Extrapolating this relation to the lower velocity dispersions of globular clusters, with $\\sigma \\sim 10 - 20$ \\kms, predicts central black holes in these objects with masses of $10^3 - 10^4 M_{\\odot}$.  Due to the small amount of gas and dust in globular clusters, the accretion efficiency of a potential black hole at the center is expected to be low. Therefore, the detection of IMBHs at the centers of globular clusters through X-ray and radio emissions is challenging \\citep{miller_2002, maccarone_2008}. Nevertheless, there is another way to detect IMBHs in globular clusters: exploring the kinematics of these systems in the central regions. This method, proposed forty years ago \\citep{bahcall_1976,wyller_1970}, has long been limited by the quality of observational datasets, since it requires velocity dispersion measurements at a spatial resolution comparable to the size of the gravitational sphere of influence for plausible IMBH masses ($1 - 2 ''$ for large Galactic globular clusters). However, with existence of the Hubble Space Telescope (HST) and with high spatial resolution ground based integral-field spectrographs, the search for IMBHs was revitalized.  \\cite{gebhardt_1997,gebhardt_2000} and \\cite{gerssen_2002} claimed the detection of a black hole of $(3.2 \\pm 2.2) \\times 10^3 M_{\\odot}$ in the globular cluster M15 from photometric and kinematic observations. After more investigations this cluster no longer appears as a strong IMBH candidate \\citep[e.g.][]{dull_1997,baumgardt_2003a,baumgardt_2005,bosch_2006}, but new detections of IMBH candidates in other clusters followed. \\cite{gebhardt_2002,gebhardt_2005} used the velocity dispersion measured from integrated light near the center of the M31 cluster G1 to argue for the presence of a $(1.8 \\pm 0.5) \\times 10^4 M_{\\odot}$ dark mass at the cluster center. The possible presence of an IMBH in G1 gained further credence with the detection of weak X-ray and radio emission from the cluster center \\citep{pooly_2006, kong_2007,ulvestad_2007}. Also, the globular cluster $\\omega$ Centauri (NGC 5139) has been proposed to host a black hole at its center \\citep{noyola_2008, noyola_2010}. The authors measured the velocity-dispersion profile with an integral field unit and used orbit based dynamical models to analyze the data. \\cite{vdm_2010} studied the same object using proper motions from HST images. They found less compelling evidence for a central black hole, but more importantly, they found a location for the center that differs from previous measurements. Both G1 and $\\omega$ Centauri have been suggested to be stripped nuclei of dwarf galaxies \\citep{freeman_1993, meylan_2001} and therefore may not be the best representatives of globular clusters. The key motivation of this work is to probe more globular clusters for the presence of IMBHs.  Further evidence for the existence of IMBHs is the discovery of ultra luminous X-ray sources at non-nuclear locations in starburst galaxies \\citep[e.g.][]{fabbiano_1989, colbert_1999, matsumoto_2001, fabiano_2001}. The brightest of these compact objects (with $L \\sim 10^{41} \\ \\mathrm{erg}\\, \\mathrm{s}^{-1}$) imply masses larger than $10^3 M_{\\odot}$ assuming accretion at the Eddington limit. Several realistic formation scenarios of black holes in globular clusters have been developed. The two main formation theories are: a) IMBHs would be Population III stellar remnants \\citep{madau_2001}, or b) they would form in a runaway merging of young stars in sufficiently dense clusters \\citep{zwart_2004, gurkan_2004, freitag_2006}. In addition \\cite{miller_2002} presented scenarios for the capture of clusters by their host galaxies and accretion in the galactic disk in order to explain the observed bright X-ray sources.  Our goal in this paper is to study the globular cluster NGC 6388. This cluster is located $11.6$ kpc away from the Sun, in the outer bulge of our Galaxy. \\cite{white_1972} assigned it a high metallicity after studying its color magnitude diagram. Later, \\cite{illingworth_1974} determined the dynamical mass of the cluster: with $\\sim 1.3 \\times 10^6 M_{\\odot}$ it belongs to the most massive clusters in the Milky Way. In addition, its high central velocity dispersion of $18.9$ \\kms  \\citep{pryor_1993}, and assuming a black-hole mass correlation with velocity dispersion, make NGC 6388 a good candidate for detecting an intermediate-mass black hole. Besides the kinematic properties, the photometric characteristics are also quite interesting. \\citet[hereafter NG06]{noyola_2006} found a shallow cusp in the central region of the surface brightness profile of NGC 6388. N-body simulations showed that this is expected for a cluster hosting an intermediate-mass black hole \\citep{baumgardt_2005}. For that reason, \\cite{lanzoni_2007} investigated the projected density profile and the central surface brightness profile with a combination of HST high-resolution and ground-based wide-field observations. They found the observed profiles are well reproduced by a multimass, isotropic, spherical King model, with an added central black hole with a mass of $\\sim 5.7 \\times 10^3 M_{\\odot}$. Also, the work of \\cite{miocchi_2007} suggests the presence of an intermediate-mass black hole in NGC 6388 as a possible explanation for the extended horizontal branch.  Another interesting fact about NGC 6388 is that it appears to contain multiple stellar populations \\cite[e.g.][]{yoon_2000,piotto_2008}. This fact makes the scenario of merging clusters plausible, leading to potentially complicated dynamics in the center. Recently, X-ray observations for NGC 6388 showed no significant signatures for an accreting IMBH  \\citep{cseh_2010}. But as already mentioned, this does not rule out a quiescent black hole. In summary, this cluster was chosen as it presents many interesting features. We measured the kinematics of the central regions, allowing us to probe the result of \\cite{lanzoni_2007} with a different method, taking the kinematic properties as well as the photometric properties into account.  The basic approach of this work is to first study the light distribution of the cluster. Photometric analysis, such as the determination of the cluster center and the measurement of a surface brightness profile, is described in section \\ref{phot}. De-projecting this profile gives an estimation of the gravitational potential produced by the visible mass. The next step is to study the dynamics of the cluster. Section \\ref{spec} gives an overlook of the FLAMES observations and data reduction and section \\ref{kin} describes the analysis of the spectroscopic data. With the resulting velocity-dispersion profile, it is possible to estimate the actual dynamical mass. The next step is to compare the data to dynamical Jeans models (section \\ref{jeans}). These models take the light profile and predict a velocity-dispersion profile, which is scaled to the data in order to obtain the mass-to-light ratio. This is done for models containing different black-hole masses until the best fit to the observed profile is found. In section \\ref{con} we summarize our results, list our conclusions and give an outlook for further studies. %  \\begin{table} \\caption{Properties of the globular cluster NGC 6388 from the references: NG=\\cite{noyola_2006}, H= \\cite{harris_1996}, M= \\cite{moretti_2009}, L=\\cite{lanzoni_2007} and PM=\\cite{pryor_1993}.}             % \\label{tab_6388}      % \\centering \\begin{tabular}{l l l} \\hline \\hline \\noalign{\\smallskip} Parameter & Value & Reference\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} RA (J200) & $17\\mathrm{h} \\ 36\\mathrm{m} \\ 17\\mathrm{s}$  & NG \\\\ DEC (J200) & $-44^{\\circ}\\ 44' \\ 08''$ & NG\\\\ Galactic Longitude l & $345.56 \u00b0$ &H\\\\ Galactic Latitude b & $-6.74 \u00b0$ &H\\\\ Distance from the Sun $R_{\\mathrm{SUN}}$ & $11.6 \\ \\mathrm{kpc}$ & M \\\\ Core Radius $r_c$ & $7.2''$ & L  \\\\ Central Concentration c & $1.8$ & L  \\\\ Heliocentric Radial Velocity V$_r$ & $81.2 \\pm 1.2 \\ \\mathrm{km}/\\mathrm{s}$ & H \\\\ Central Velocity Dispersion $\\sigma$ & $18.9 \\ \\mathrm{km}/\\mathrm{s}$ & PM \\\\ Age & $ (11.5 \\pm 1.5) \\ \\mathrm{Gyr}$ & M \\\\ Metallicity $[\\mathrm{Fe}/\\mathrm{H}]$ & $-0.6 \\ \\mathrm{dex}$ & H \\\\ Integrated Spectral Type  & G2 &H \\\\ Reddening E(B-V) & $0.38 $ & M \\\\ Absolute Visual Magnitude $M_{Vt}$ & $-9.42 \\ \\mathrm{mag}$ & H  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\label{con}   We derived the mass of a potential intermediate-mass black hole at the center of the globular cluster NGC 6388 by analyzing spectroscopic and photometric data. With a set of HST images, the photometric center of the cluster was redetermined and the result of NG06 confirmed. Furthermore, a color magnitude diagram as well as a surface brightness profile, built from a combination of star counts and integrated light, were produced. The spectra from the ground-based integral-field unit ARGUS were reduced and analyzed in order to create a velocity map and a velocity-dispersion profile. In the velocity map, we found signatures of rotation or at least complex dynamics in the inner three arcseconds (0.15 pc) of the cluster. We derived a velocity-dispersion profile by summing all spectra into radial bins and applying a penalized pixel fitting method.  Using the surface brightness profile as an input for spherical Jeans equations, a model velocity-dispersion profile was obtained. We ran several isotropic models with different black-hole masses and scaled them to the observed data in order to measure the mass-to-light ratio. Using $\\chi^2$ statistics, we were able to find the model which fits the observed data best. We run Monte Carlo simulations on the inner points of the surface brightness profile in order to estimate the errors resulting from the particular choice of a light profile. From this, we determined the black-hole mass and the mass-to-light ratio as well as the scatter of these two results. The final results, with $68 \\%$ confidence limits are: A black-hole mass of $ (17 \\pm 9) \\times 10^3 M_{\\odot}$ and a global mass-to-light ratio of $M/L_V = (1.6 \\pm 0.3) \\ M_{\\odot}/L_{\\odot}$. In addition, we run anisotropic models. The confidence of black hole detection is decreasing rapidly with increasing radial anisotropy. However, using N-body simulations, we found that in a relaxed cluster such as NGC 6388 an anisotropy higher than $\\beta=0.2$ inside the half mass radius will be erased within a few relaxation times.  With respect to the black-hole mass estimate of $5.7 \\times 10^3 \\ M_{\\odot}$ by \\cite{lanzoni_2007}, which they derived from photometry alone, our derived black-hole mass is larger by a factor of three with our surface brightness profile profile. Using their light profile, we obtained a black-hole mass of $M_{\\bullet} = (11 \\pm 5) \\times 10^3 M_{\\odot}$ which includes their value within one sigma error. \\cite{cseh_2010} obtained deep radio observations of the inner regions of NGC 6388 and discussed  different accretion scenarios for a possible black hole. They found an upper limit for the black-hole mass of $\\sim 1500 M_{\\odot}$ since no radio source was detected at the location of any Chandra X-ray sources. However, inserting our black-hole mass in equation 5 of \\cite{cseh_2010} and using their assumptions as well as the Chandra X-ray luminosity of the central region, results in minimal but possible accretion rates and conversion efficiency ($\\varepsilon \\eta \\sim 10^{-7}$). Considering the fact that supermassive black holes have masses not much higher than $0.2$ \\% of the mass of their host systems \\citep{marconi_2003,haering_2004}, our mass with $0.9 \\%$ seems to be higher than the predictions for larger systems. Globular clusters in contrast lose much of their mass during their evolution. This could naturally result in higher values of black-hole mass - host system mass ratios. Due to the complicated dynamics, the $1 \\sigma$ uncertainties are 40 \\% for our black-hole mass and 10 \\% for the derived $M/L_V$. Figure \\ref{m_sigma} shows the position of NGC 6388 in the black-hole mass velocity dispersion relation. With our derived effective velocity dispersion of $\\sigma_e = (18.9 \\pm 0.3)$ \\kms the results for NGC 6388 seem to coincide with the prediction made by the $M_{\\bullet} - \\sigma$ relation.  The mass-to-light ratio of $M/L_V = (1.7 \\pm 0.3) \\, M_{\\odot}/L_{\\odot}$ derived in this work is consistent within the error bars with the dynamical results of \\citet[$M/L_V\\sim 1.8 $]{McLaughlin_2005} but slightly lower than the value predicted by stellar population models. According also to \\cite{McLaughlin_2005}, we would expect a value of $M/L_V = (2.55 \\pm 0.28) \\, M_{\\odot}/L_{\\odot}$ for $\\mathrm{[Fe/H]}=-0.6 \\,  \\mathrm{dex}$ and  $(13 \\pm 2) \\, \\mathrm{Gyr}$. \\citet{baumgardt_2003} have shown that dynamical evolution of star clusters causes a depletion of low-mass stars from the cluster. Since these contribute little light, the $M/L_V$ value drops as the cluster evolves. Dynamical evolution could therefore explain part of the discrepancy between theoretical and observed $M/L_V$ values.  Summarizing, it can be said that the globular cluster NGC 6388 shows a variety of interesting features in its photometric properties (e.g. the extended horizontal branch) as well as in its kinematic properties (e.g. the high central velocity dispersion or the rotation like signature in the velocity map). In this work we investigated the possibility of the existence of an intermediate-mass black hole at the center. With simple analytical models we were already able to reproduce the shape of the observed velocity-dispersion profile very well if the cluster hosts an IMBH. Future work would be to compute detailed N-body simulations of NGC 6388 in order to verify whether a dark cluster of remnants at the center would be an alternative. Additionally, it is crucial to obtain kinematic data at larger radii to constrain the mass-to-light ratio and the models further out. Also proper motions for the central regions would further constrain the black hole hypothesis.  The study of black holes in globular clusters is currently limited to a handful of studies. For this reason, it is necessary to probe a large sample of globular clusters for central massive objects in order to move the field of IMBHs from ``whether\" to ``under which circumstances do\" globular clusters host them. Thus tying this field to the one of nuclear clusters and super-massive black holes. From the experience we gathered with NGC 6388, we can say that the dynamics of globular clusters is not as simple as commonly assumed and by searching for intermediate-mass black holes, we will also be able to get a deeper insight into the dynamics of globular clusters in general.   \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{M_sigma.eps} \\caption{The $M_{\\bullet} - \\sigma$ relation for galaxies at the high mass range (filled circles) and the first globular clusters with potential black hole detection (filled stars). The slope of the line $\\log (M_{\\bullet}/M_{\\odot}) = \\alpha + \\beta \\, \\log ( \\sigma / 200 \\mathrm{km\\, s^{-1}\\, })$ with $(\\alpha, \\beta) = (8.12 \\pm 0.08, 4.24 \\pm 0.41)$ was taken from \\cite{gultekin_2009}. The mass of the black holes of $\\omega$ Centauri, G1 and the upper limits of M15 and 47 Tuc were obtained by \\cite{noyola_2010}, \\cite{gebhardt_2005}, \\cite{bosch_2006} and \\cite{47tuc}, respectively. Overplotted is the result of NGC 6388 with our derived effective velocity dispersion $\\sigma_e = 18.9$ \\kms.} \\label{m_sigma} \\end{figure}"
    },
    "1107/1107.5560.txt": {
        "abstract": "Recently a new particle physics model called Bound Dark Matter (BDM) has been proposed \\cite{delaMacorra:2009yb} in which dark matter (DM) particles are massless above a threshold energy ($E_c$) and acquire mass below it due to nonperturbative methods. Therefore, the BDM model describes DM particles which are relativistic, hot dark matter (HDM) in the inner regions of galaxies and describes nonrelativistic, cold dark matter (CDM) where halo density is below $\\rho_c\\equiv E_c^4$. To realize this idea in galaxies we use a particular DM cored profile that contains three parameters: a typical scale length ($r_s$) and density ($\\rho_0$) of the halo, and a core radius ($r_c$) stemming from the relativistic nature of the BDM model. We test this model by fitting rotation curves of seventeen Low Surface Brightness (LSB) galaxies from The HI Nearby Galaxy Survey (THINGS). Since the energy $E_c$ parameterizes the phase transition due to the underlying particle physics model, it is independent on the details of galaxy and/or structure formation and therefore the DM profile parameters $r_s, r_c, E_c$ are constrained, leaving only two free parameters. The high spatial and velocity resolution of this sample allows to derive the model parameters through the numerical implementation of the $\\chi^2$-goodness-of-fit test to the mass models. We compare the fittings with those of  Navarro-Frenk-White (NFW), Burkert, and Pseudo-Isothermal (ISO) profiles. Through the results we conclude that the BDM profile fits better, or equally well, than NFW, Burkert, and ISO profiles and agree with previous results implying that cored profiles are preferred over the N-body motivated cuspy profiles.  We also compute 2D likelihoods of the BDM parameters $r_c$ and $E_c$ for the different galaxies and matter contents, and find an average galaxy core radius $r_c = 300$ pc and a transition energy between hot and cold dark matter at $E_c = 0.11^{+0.21}_{-0.07} {\\rm \\ eV}$ when the DM halo is the only component, therefore the maximum dark matter contribution in galaxies. In a more realistic analysis, as in Kroupa mass model, we obtain a core $r_c = 1.48$ kpc, and energy $E_c = 0.06^{+0.07}_{-0.03} {\\rm \\ eV}$. ",
        "introduction": "\\label{sec:introduction}  In the last decades intense work to understand the distribution of dark matter (DM) in galaxies has been published \\cite{KrAl78,Tr87,SoRu01, deBlok10}. Among the considered, late-type Low Surface Brightness (LSB) galaxies are of special interest since it is believed they are dominated by DM, and high resolution $HI$, $H_{\\alpha}$ and optical data can help to distinguish among the different DM profiles proposed in the literature.  There are essentially two types of profiles, the ones stemming from cosmological $N$-body simulations that have a cusp in its inner region, e.g. Navarro-Frenk-White (NFW) profile \\cite{Navarro:1995iw,Navarro:1996gj}. On the other hand, the phenomenological motived cored profiles, such as the Burkert or Pseudo-Isothermal (ISO) profiles \\cite{Bu95}. Cuspy and cored profiles can both be fitted to most LSB rotation curves, but with a marked preference for a cored inner region with constant density. Furthermore, cuspy profiles that do fit to galaxies in many cases suffer from a parameter inconsistency, since their concentrations ($c$) are too low and velocities $V_{200}$ are too high in comparison to the ones expected from cosmological simulations \\cite{KuMcBl08,deBlok10}. Some other fits indicate that cuspy profiles require to tune fine the observer's line of sight with axisymmetric potentials \\cite{KuMcMi09}. However, different systematics may play an important role in the observations such as noncircular motions, resolution of data and other issues \\cite{Swater99,vandenBosch:1999ka,Swaters:2000nt, Rhee:2003vw,Simon:2003xu,deBlok:2008wp}. And there are attempts to reconcile both approaches through evolution of DM halo profiles including baryonic processes  \\cite{SpGiHa05,Oh:2010mc} to transform cuspy to shallower profiles that follow the solid-body velocity curve ($v \\sim r$) observed in late-type LSB. However, fthis issue is a matter of recent debate, see e.g. \\cite{Pontzen:2011,Ogiya:2011ta} for a recent discussion.  Alternatively to the above models, recently one of us proposed a profile based on a particle physics model called bound dark matter (BDM) \\cite{delaMacorra:2009yb}, in which DM particles at high densities, as in galactic inner regions, are relativistic, i.e. Hot DM (HDM), but in the outer regions they behave as standard CDM. To realize this idea in galaxies we use in the present work a particular DM cored profile that contains three parameters: a typical scale length ($r_s$) and density ($\\rho_0$) of the halo and a core radius ($r_c$) stemming from the relativistic nature of the BDM model. The galactic core density is given by $\\rho_c\\equiv E_c^4$ and the profile properties will determine the energy scale of the particle physics model. We will show  that the LSB rotation curves yield a phase transition energy scale $E_c$, between HDM and CDM for our BDM profile, at $E_c = 0.11^{+0.21}_{-0.07} {\\rm \\ eV}$, when we consider only DM and $E_c = 0.06^{+0.07}_{-0.03}$ when considering gas, DM halo, and the minimum contribution of stars \\cite{delaMacorra:2011df}. The $E_c$ parameter is a new fundamental scale  for  DM which can be theoretically determined using  gauge group dynamics, i.e. it does not depend on the properties of galaxies, once the gauge group is known. However, even though we propose $E_c$  as a new fundamental constant for DM, close related to the mass of our DM particle, and  it is important to note that its value is not yet known and we require observational evidence to determine  it.  The same is true for all masses of the standard model of particles, i.e. their value is not predicted by the standard model and it is the experimental results that fixed them in a consistent manner. We use here the information on galactic rotation curves to determine $E_c$. However, the extracting of $E_c$  does depend not only on the quality of the observational data, the mass models used but also on the choice of DM profile.   To perform this task, we use The HI Nearby Galaxy Survey (THINGS), which collects high spectral resolution data revealing extended measurements of gas rotation velocities and circular baryonic matter trajectories \\cite{Walter:2008wy}. Given these properties it is adequate to test the above-mentioned DM profiles with THINGS, which has been used to test  different core/cusp mass profiles. For disk-dominated galaxies the core and cusp profiles fit equally well, however for LSB galaxies there is a clear preference for core profiles over the cuspy models \\cite{deBlok:2008wp}. Analysis of different high resolution datasets have confirmed this tendency in past recent years \\cite{SpGiHa05,KuMcBl08,Salucci:2010pz,Gentile:2004tb,Salucci:2007tm}.  In this work, we present a study of the rotation curves using BDM, NFW, Burkert and ISO DM profiles, taking into account the contribution from different mass models: i) DM alone; ii) DM and gas; and iii) DM, gas and the stellar disk.  By fitting the models to the data we find that the BDM profile fits equally well or better than the cored profiles and only for a few LSB galaxies our model resembles that of NFW.  The kinematics of stars bring a very challenging problem in the analysis, mainly due to the uncertainty of the mass-to-light ratio ($\\gs$) and to its dominant behavior close to the galactic center. Some considerations have been made in order to reduce this uncertainty in the parameters \\cite{vanAlbada,Salpeter:1955it,Kroupa:2000iv,Bottema:1997qe}, but still the stellar contribution is not well known. In our case, we study four different $\\gs$ models and the most inner part of the data enable us to compute the BDM parameters $r_c$ and $E_c$, both related to the galactic core.  We present the 1$\\sigma$ and $2\\sigma$ likelihood  contour plots of these parameters for the different galaxies and mass models. We notice that being $r_c$ different for each mass model it is consistent within the $2\\sigma$ error for each galaxy. Additionally, we find that the galaxy core radius is the order of $r_c \\sim 300 {\\rm \\ pc}$ and the $E_c \\sim \\mathcal{O}(0.1 {\\rm \\ eV})$ which gives the transition between hot and cold DM, this is shown in preliminary study \\cite{delaMacorra:2011df}. It is also interesting to note the coincidence in the magnitude of the sum of neutrino masses with the magnitude of $E_c$ obtained in this analysis, and how this  could  open an interesting  connection between the generation of DM and neutrinos masses, but this will be presented elsewhere \\cite{NBDM}. The magnitude of  the energy transition $E_c$  is of the same order as the upper limit of the total mass of the neutrinos $\\sum m_\\nu < 0.58 {\\rm \\ eV} (95\\% {\\rm \\ CL)}$ from the seven year Wilkinson Microwave Anisotropic Probe (WMAP), Baryon Acoustic Oscillation (BAO), and the Hubble constant ($H_0$) \\cite{Larson:2010gs,Komatsu:2010fb}. In addition, the recent analysis of the Atacama Cosmology Telescope \\cite{Dunkley:2010ge} that combined CMB data  with measurements of BAO and the Hubble constant reported an excess of the effective number of relativistic degrees of freedom, $N_{eff} = 4.6 \\pm 0.8$, in consistency with Ref. \\cite{Komatsu:2010fb}, opening the possibility to have extra nonstandard-model relativistic degrees of freedom, as our BDM particle.  We organized this work as follows:  In Sec. \\ref{SPM} we explain the particle model behind the BDM profile. In Sec. \\ref{modelo} we  present our BDM profile, which has NFW as a limit at low energies. In Sec. \\ref{sample} we describe the galaxy sample considered for the rotation curves study. The different mass models and components (gas, stars, DM halo) are presented  in detail in Sec. \\ref{MaMo}, while the different hypothesis of the mass-to-light ratio are discussed in subsection \\ref{MaMoEsDi}. The results and conclusions are presented in Sec. \\ref{results} and \\ref{conclusion}, respectively. For convenience, some figures and tables are shown in the Appendices. We show in Appendices \\ref{apend: diet-salpeter} and \\ref{apendix:iso_burkert_result} the fitting values corresponding to the free parameters for some of the cored profiles for the maximum stellar contribution. Finally, in Appendix \\ref{apend:LikelihoodPlot} we present the rotation curves of the different profiles and  galaxies and we include the confidence contours levels, pointing out the considerations made for each galaxy. ",
        "conclusions": "We have presented the analysis of the rotation curves using four different DM profiles (BDM, NFW, Burkert, and ISO) and five different disk mass components (only DM (min.disk), DM plus gas (min.disk+gas), DM plus gas plus stellar disk (Kroupa, diet-Salpeter, Free $\\gs$) for a subsample of the THINGS galaxies as described in Sec. \\ref{MaMo}.  The new tested profile, BDM, is introduced in Sec. \\ref{modelo} and  corresponds to CDM particles below certain critical energy density $\\rc = E_c^4$ while it behaves as HDM above $\\rc$ with a core radius $r_c$. Therefore, the resulting profile behaves as NFW's for large radius, but it has a core in the central region.  We have found that fits to the THINGS galaxies favor core profiles over the cuspy NFW in accordance to general results found in other works \\cite{deBlok:2008wp,deBlok10} for the same galaxies. We studied the cases when the five different disk mass components are taking into account and conclude that BDM is in good agreement with observations for all the disk mass scenarios and fits better or equally well than Burkert and ISO profiles.  For three galaxies (Group C) the fits to the BDM profile demanded $r_c/r_s <10^{-6}$ and in these cases the $\\chi^2$ are almost identical to NFW's.  For these galaxies a cuspy profile fits best because there is not enough observational data close to the galactic center to test it.  The observational resolution of the THINGS sample is of high quality, but one still needs data closer to the center of some galaxies, on scales smaller than 200 {\\rm pc}, in order to discern between cored or cuspy profiles.  This is because stars pose a very challenging problem when testing the core-cusp problem, largely due to the uncertainty of the mass-to-light ratio and their dominant behavior close to the center of the galaxy. We have performed an analysis of the inner regions of  Group B and C galaxies, for which $r_c \\leq r_s$, extracting information from the rotation curves consistent with a core profile.  Confidence level contours show that the core radius depends on the number of disk mass components taking into account in the analysis.  We found that $r_c$ is highly constrained if the stellar disk has a dominant behavior close to the center of the galaxy.  We computed the 1 and 2$\\sigma$ confidence levels for the BDM parameters $r_c$ and $E_c$ for the different galaxies and mass models. We found that the energy of transition between HDM and CDM takes places for most galaxies with an energy $E_c = \\mathcal{O}(0.1{\\rm \\, {\\rm eV}})$ and most of the galaxies have a core radius $r_c \\sim \\mathcal{O}(300 \\,$ {\\rm pc}), see Table \\ref{tab:statistics_BDM}. We found that $r_c$ is highly constrained if the stellar disk has a dominant behavior close to the center of the galaxy. However, we notice that even when $r_c$ is different for each disk  mass model there is an interval of values where $r_c$ is consistent for all mass models within the $1\\sigma$ for most galaxies and some within the $2\\sigma$ error, as shown in Figures \\ref{tab:conflevel_A}, \\ref{tab:conflevel_B}, \\ref{tab:conflevel_C}, and \\ref{tab:conflevel_inner B}.  This provides us with enough tools to later test the Universe background evolution and large scale structure formation in order to continue testing the particle-physics-motivated BDM model. This energy transition is similar to the mean energy of a relativistic neutrino (see e.g. \\cite{Larson:2010gs}) which is such that $\\sum m_\\nu < 0.58 {\\rm \\ eV} (95\\% {\\rm \\ CL)}$.  The dispersion on $E_c$ is much smaller than that of $r_c$ and this is consistent with our BDM model since $E_c$ is a new fundamental scale for DM model while $r_c$ depends on $r_s,\\rho_0$ which depend on the formation and initial conditions for each galaxy. We test the BDM profile with a fixed energy $E_c$ finding that the proposed profile still being competitive compare with NFW profile. BDM profile has a rich -particle physics motivated- structure that, through its transition from CDM to HDM (when $\\rho \\simeq \\rho_c$),  allows for a more physical explanation of the rotation curves of the different galaxies treated here."
    },
    "1107/1107.5548.txt": {
        "abstract": "{We constrain the rotation angle $\\alpha$ of the linear polarization of CMB photons using  the large angular scale (up to $\\sim 4^\\circ$) signal in WMAP 7 year data. At these scales, the CMB polarization pattern probes mostly the reionization era. A genuine rotation may be interpreted as cosmological birefringence, which is a well known tracer of new physics, through the breakdown of fundamental symmetries. Our analysis provides $\\alpha = -1.6^{\\circ} \\pm 1.7^{\\circ} \\, (3.4^{\\circ})$ at $68\\%$ ($95\\%$) C.L. for the multipoles range $\\Delta \\ell = 2-47$ not including an instrumental systematic uncertainty that the WMAP team estimates at $\\pm 1.5^\\circ$. This improves the bound obtained by WMAP team (Komatsu et al., 2010). Moreover we show, for the first time at low multipoles, the angular power spectrum $\\alpha_\\ell$ in search of a possible scale dependence of the birefringence effect. Our findings are compatible with no detection at all angular scales probed here. We finally forecast the capabilities of Planck in tightening the present constraints.} ",
        "introduction": "\\label{intro}  The Cosmic Microwave Background (CMB) can be used to constrain fundamental symmetry breaking. Under the Standard Model of particle physics, the CMB temperature to $B$ mode polarization correlation ($TB$) and the $E$ mode to $B$ mode polarization correlation ($EB$) are expected to vanish since parity is conserved \\cite{Lue:1998mq,saito}. The latter prediction follows from the purely electromagnetic, charge blind nature of the physics of CMB anisotropies \\cite{carroll90,carroll91} and becomes observable due to the opposite handedness of $(T,E)$ vs $B$ modes \\cite{kks-seljak}. A non null measurement of either $TB$ or $EB$ would thus hint at new physics beyond the Standard Model (or suggest a careful reanalysis of the dataset in search for un-removed systematic effects).  A number of models predict non-standard parity (P) and CP violations (where C stands for charge conjugation), as well as CPT violations (T standing for time reversal) and the associated breakdown of Lorentz invariance \\cite{greenberg}. Parity violations in the electromagnetic sector are associated with an \\textit{in vacuo} rotation of the linear polarization direction of a photon \\cite{carroll90,carroll91}, an effect known as cosmological birefringence. Several tests have been set forth, either in earthly and orbital laboratories \\cite{bluhm-mewes} or through cosmological observations \\cite{carroll90,carroll91,amelino,Alighieri:2010pu}. These violations should also impact the CMB, whose statistical properties are derived under the assumption of symmetry conservation. The CMB is a powerful probe for such analyses for two main reasons. First, it is generated in the early universe, when the physics at the stake was not obviously identical to present. Secondly, the long look-back time of CMB photons may render tiny violations to the electromagnetic Lagrangian observable, since such effects usually accumulate during propagation.  CMB polarization arises at two distinct cosmological times: the recombination epoch ($z \\simeq 1100$) and the reionization era at $z \\lesssim 11$ \\cite{dodelson}. When the CMB field is expanded in spherical harmonics, the first signal mostly shows up at high multipoles, since polarization is generated through a causal process and the Hubble horizon at last scattering only subtends a degree sized angle. The later reionization of the cosmic fluid at lower redshift impacts the low $\\ell$ instead. These two regimes need to be taken into account when probing for cosmological birefringence, since they can be ascribed to different epochs and, hence, physical conditions.  %For instance, the presence of a primordial homogeneous %\\cite{scafer} or helical \\cite{pogosian} magnetic field  would induce %Faraday rotation and non-zero $TB$ correlations. Parity-asymmetric %gravity dynamics during inflation could cause unbalance in %left and right-handed gravitational waves, which impacts $TB$ and $EB$ \\cite{Lue:1998mq,saito}. In general, models in high energy physics with non-standard %parity-violating interactions also predict $TB$ and $EB$ signals different from zero \\cite{maity}. %A popular model for which parity is broken in the photon sector is the Chern-Simons perturbation to the Maxwell Lagrangian \\cite{carroll90}: %\\[ %  \\Delta{\\mathcal L}= -\\frac{1}{4}\\,  p_{\\mu} %  \\epsilon^{\\mu\\nu\\rho\\sigma}F_{\\rho\\sigma} A_{\\nu}\\;, %\\] %where ${F}^{\\mu\\nu}$ is the Maxwell tensor and $A^\\mu$ the four-potential. %The four-vector $p_{\\mu}$ can be interpreted in several ways, e.g., the derivative of the quintessence field or the gradient of a function of the Ricci scalar %\\cite{Davoudiasl}. In any case a P violation always arises provided %that the timelike component of $p_{\\mu}$ does not vanish. C and T symmetries are then unbroken %so CP and CPT are not conserved. Given that $p^\\mu$ selects a %preferred direction in spacetime, Lorenz invariance cannot be preserved.  Historically, the effect has been first constrained by measuring polarized light from high redshift radio galaxies and quasars \\cite{carroll90,carroll91,Alighieri:2010pu,carroll97,Nodland,Eisenstein,Leahy}, see \\cite{Alighieri:2010eu} for a new analysis on ultra-violet polarization of distant radio galaxies. For recent, polarization oriented CMB observations \\cite{Pryke:2008xp,Komatsu:2010fb,Piacentini:2005yq,Montroy:2005yx} it has become feasible to measure $TB$ and $EB$ correlations, other than $TT$, $TE$ and $EE$ correlations. While no detection has been claimed to date, polarization data have been used to derive constraints on the birefringence angle \\cite{Komatsu:2010fb,Feng:2006dp,Cabella:2007br,Wu:2008qb}. %The tight control of  systematic effects is essential. See \\cite{Pagano:2009kj} for an analysis that takes polarimeter calibration into account. %However in the current analysis we do not address the issue of the systematic effects whose detailed analysis goes beyond the scope of the paper. %Our new results are based on the WMAP data and we do rely on their original estimate for the uncertainty coming from systematic effects on the birefringence angle \\cite{Komatsu:2010fb}. %For what concerns the Planck forecast we provide, we cannot give an estimate on the errors coming from Planck detectors since those are not yet publicly available. %However, following  \\cite{Pagano:2009kj}, we will provide an estimate of the level of accuracy that polarimeter calibration has to satisfy in order to make the error %coming from this systematic effect smaller than the statistical one.  The goal of this paper is to constrain the cosmological birefringence angle at large angular scale considering the latest WMAP data release. Our method naturally provides also a scale dependence for such an angle. %This is done adopting a method which is capable also to built the spectrum for such an angle (see below).  In the current analysis we do not address the issue of the systematic effects whose detailed analysis goes beyond the scope of the paper. Our new results are based on the WMAP data and we do rely on their original estimate of the uncertainty on the birefringence angle coming from systematic effects \\cite{Komatsu:2010fb}. For what concerns the Planck forecast we provide, we cannot give an estimate on the errors coming from Planck detectors since those are not yet publicly available. However, following  \\cite{Pagano:2009kj}, we will provide an estimate of the level of accuracy that polarimeter calibration has to satisfy in order to make the error coming from this systematic effect smaller than the statistical one.  In the limit of constant birefringence angle, $\\alpha$, the angular power spectra of CMB anisotropies, assuming $C_{\\ell}^{TB} = C_{\\ell}^{EB} = 0 $, are given by \\cite{Lue:1998mq,Feng:2006dp,Feng:2004mq,Komatsu:2008hk} \\footnote{See \\cite{Galaverni:2009zz,Finelli:2008jv} as an example of computation that takes into account the time dependence of $\\alpha$ in a specific model of pseudoscalar fields coupled to photons. See \\cite{Li:2008tma,Kamionkowski:2010ss,Gubitosi:2011ue} as examples of non-isotropic birefringence effect.},  %The polarization rotation can be parametrized by the angle $\\alpha$, namely the birefringence angle,  that, in the limit of constant $\\alpha$ %\\footnote{See \\cite{Galaverni:2009zz,Finelli:2008jv} as an example of computation that takes into account the time dependence of $\\alpha$ in a specific model of pseudoscalar fields coupled to photons. %See \\cite{Li:2008tma,Kamionkowski:2010ss,Gubitosi:2011ue} as examples of non-isotropic birefringence effect.}, %%, i.e. the rotation of the CMB polarization direction depends in a peculiar way on the observation direction.}) %impacts the angular power spectra %of CMB anisotropies as follows  \\cite{Lue:1998mq,Feng:2006dp,Feng:2004mq,Komatsu:2008hk} %\\begin{itemize} %\\item CMB anisotropies %\\item Testing CPT symmetry and introduction of the Birefringence angle $\\alpha$, \\cite{Lue:1998mq,Feng:2004mq,Komatsu:2008hk,Feng:2006dp} \\begin{eqnarray} \\langle C_{\\ell}^{TE,obs} \\rangle &=& \\langle C_{\\ell}^{TE} \\rangle \\cos (2 \\alpha) \\, , \\label{TEobs} \\\\ \\langle C_{\\ell}^{TB,obs} \\rangle &=& \\langle C_{\\ell}^{TE} \\rangle \\sin (2 \\alpha) \\, , \\label{TBobs} \\\\ \\langle C_{\\ell}^{EE, obs} \\rangle &=& \\langle C_{\\ell}^{EE} \\rangle \\cos^2 (2 \\alpha) + \\langle C_{\\ell}^{BB} \\rangle \\sin^2 (2 \\alpha) \\, , \\label{EEobs} \\\\ \\langle C_{\\ell}^{BB, obs} \\rangle &=& \\langle C_{\\ell}^{BB} \\rangle \\cos^2 (2 \\alpha) + \\langle C_{\\ell}^{EE} \\rangle \\sin^2 (2 \\alpha) \\, , \\label{BBobs} \\\\ \\langle C_{\\ell}^{EB, obs} \\rangle &=& {1 \\over 2} \\left( \\langle C_{\\ell}^{EE} \\rangle + \\langle C_{\\ell}^{BB} \\rangle \\right) \\sin (4 \\alpha) \\, . \\label{EBobs} \\end{eqnarray} These equations follow from \\begin{eqnarray} a_{\\ell m}^{T,obs} &=& a_{\\ell m}^{T} + a_{\\ell m}^{n,T} \\, , \\label{almT} \\\\ a_{\\ell m}^{E,obs} &=& a_{\\ell m}^{E} \\cos (2 \\alpha) +  a_{\\ell m}^{B} \\sin (2 \\alpha) + a_{\\ell m}^{n,E} \\, , \\label{almE}\\\\ a_{\\ell m}^{B,obs} &=& a_{\\ell m}^{B} \\cos (2 \\alpha) +  a_{\\ell m}^{E} \\sin (2 \\alpha) + a_{\\ell m}^{n,B} \\, , \\label{almB} \\end{eqnarray} where we have explicitly written the contribution due to residual instrumental noise present on the maps. As discussed above, the amount of rotation depends on the effective travel time of the photon and the effect is scale (and, hence, harmonic multipole $\\ell$) dependent \\cite{liu}. In the present paper, we focus on the low $\\ell$ contribution, with a twofold aim: from the one side, we provide a new constraint on $\\alpha$ for low $\\ell$ using, for the first time (to our knowledge) a linear estimator in combination with an optimal angular power spectrum estimator (APS) for such a regime \\footnote{We stress here that, strictly speaking, it is the adopted APS estimator to be optimal. In particular we have considered the QML method which is called optimal  \\cite{Tegmark:1996qt,Tegmark:2001zv} since it is unbiased and minimum variance (it provides the smallest error bars provided by the Fisher-Cramer-Rao inequality). See also Subsection \\ref{simulations} and \\ref{datset} and reference therein.}; from the other side, we give the first measurement of  the spectrum of $\\alpha$ at low $\\ell$. Note that the technique used by WMAP team does not allow for possible scale dependence of $\\alpha$.  Our constraint reads $\\alpha = -1.6^{\\circ} \\pm 1.7^{\\circ} \\, (3.4^{\\circ})$ at $68\\%$ ($95\\%$) C.L. for $\\Delta \\ell = 2-47$. Considering $\\Delta \\ell = 2-23$ we obtain $\\alpha = -3.0^{\\circ +2.6^{\\circ}}_{\\phantom{a}-2.5^{\\circ}}$ at $68\\%$ C.L. and $\\alpha = -3.0^{\\circ +6.9^{\\circ}}_{\\phantom{a}-4.7^{\\circ}}$ at $95\\%$ C.L.. This is the same multipole range considered by the WMAP team in \\cite{Komatsu:2010fb} (the only other result available in the literature at these large angular scales) where  with a pixel based likelihood analysis they obtain $\\alpha^{{\\rm WMAP} \\, 7 yr} = -3.8^{\\circ} \\pm 5.2^{\\circ}$ at $68\\%$ C.L..  %The plan of this paper is as follows... %\\begin{itemize} %\\item How $\\alpha$ has been constrained in literature \\cite{Alighieri:2010pu,Galaverni:2009zz,Finelli:2008jv,Alighieri:2010eu}. %\\item Observe that up to now, (1) the spectrum $\\alpha$ vs $\\ell$ has been provided only at high $\\ell$ and (2) at low $\\ell$ the technique used by WMAP team do not allow for the building of such a spectrum %\\item Aim of the present paper is twofold: 1) provide the spectrum of $\\alpha$ at low $\\ell$ and (2) %give a new constraint on $\\alpha$ for low $\\ell$ using an estimator that has not been used in such a regime so far in literature %\\item Distinction of Physics of CMB photons for low and high $\\ell$ %\\item WMAP 7 year results and Planck forecast %\\end{itemize}  The paper is organized as follows: in Section \\ref{methodology} we present the methodology adopted and describe the  simulations performed and the tests proposed; %with a short description of the considered estimators; in Section \\ref{results} we show the results obtained (i.e.\\ a new and more stringent constraint of $\\alpha$ and its spectrum at low $\\ell$). In Section \\ref{conclusions} we draw our main conclusions. In Appendix \\ref{montecarlo} further details about the behavior of the considered estimators are given. %An alternative approach is shown in Appendix \\ref{frequentist}. ",
        "conclusions": "\\label{conclusions} We have performed an analysis of the WMAP 7 low resolution data in search for a birefringence effect in the CMB. The large angular scale considered here probes the reionization era, as opposed to other higher resolution analyses that are more sensitive to last scattering epoch. The combination of an optimal estimator for the CMB angular power spectra and a linear estimator, Eq. (\\ref{DTB}) and (\\ref{DEB}), provides constraints in the range $\\Delta \\ell= 2 - 23$ which are tighter by a factor $\\sim 2$ than WMAP's  own analysis, the only other results available in the literature at these large angular scales. We report our best constraint $\\alpha = -1.6^{\\circ} \\pm 1.7^{\\circ} \\, (3.4^{\\circ})$ at $68\\%$ ($95\\%$) C.L. considering $\\Delta \\ell= 2 - 47$. We also provide, for the first time here, an angular spectrum of the birefringence angle $\\alpha$ over the multipole range $\\ell \\lesssim 50$, detecting no significant deviation from the hypothesis of a scale independent $\\alpha$. We have, finally, forecasted the improvement expected from Planck data, finding that it should outnumber WMAP by a factor above $16$ in terms of standard deviations."
    },
    "1107/1107.2114_arXiv.txt": {
        "abstract": " ",
        "introduction": "} Star clusters are popular tracers of the star formation process, both on stellar and galactic scales \\citep[e.g.][]{smith07,pflamm07,bastian10}. This approach assumes a certain fraction of star formation that produces bound stellar clusters, also dubbed the {\\it cluster formation efficiency} (CFE). It has recently become clear that the CFE is not the product of two different `modes' of dispersed and clustered star formation. Instead, star formation proceeds according to a continuous density spectrum, of which the high-density end yields bound star clusters and the low-density end consists of unbound associations \\citep{bressert10,gieles11}. This suggests the existence of a (possibly varying) critical density that allows for the formation of bound structure. If the density spectrum of star formation would vary with environment, this would also imply that the CFE is environmentally dependent -- dense star-forming regions should then yield a high CFE (see e.g. \\citealt{adamo11} for a discussion).  It was originally thought that {\\it infant mortality} (disruption by gas expulsion) would be the key mechanism behind a CFE lower than 100\\% \\citep{lada03}. The question thus arises how this new picture, in which stars do not form in clusters but in a hierarchical setting, connects to the concept of infant mortality, and also which other (environmental) processes may determine whether stellar structure ends up being bound. The following sections address the different mechanisms at play, and how Gaia can be used to estimate their relative contributions to the CFE. ",
        "conclusions": ""
    },
    "1107/1107.4504_arXiv.txt": {
        "abstract": "Analyzing the imprint of  relic gravitational waves (RGWs) on the cosmic microwave background (CMB) power spectra provides a way to determine the signal of RGWs. In this Letter, we discuss a statistical bias, which could exist in the data analysis and has the tendency to overlook the RGWs. We also explain why this bias exists, and how to avoid it. ",
        "introduction": "}  A stochastic background of relic gravitational waves was produced in the very early stage of the Universe due to the superadiabatic amplification of zero point quantum fluctuations of the gravitational field \\cite{grishchuk,star}. The relic gravitational waves have a wide range of spreading of the spectra, and their detection provides a direct way to study the physics in the early Universe.  Recently, there have been several experimental efforts to constrain the amplitude of relic gravitational waves in different frequencies. Among various direct observations, LIGO S5 has experimentally obtained so far the most stringent bound $\\Omega_{\\rm gw}(f)\\le6.9\\times10^{-6}$ around $f\\sim 100$Hz \\cite{ligo}, which will be much improved by future observations, including the third-generation Einstein Telescope \\cite{et}. The timing studies on the millisecond pulsars by the PPTA and EPTA teams also reported upper limits $\\Omega_{\\rm gw}(f)\\lesssim 10^{-8}$ at $f\\sim 1/{\\rm yr}$ \\cite{ppta,epta}. In addition, there are two bounds on the integration $\\int \\Omega_{\\rm gw}(f)d\\ln f\\lesssim 1.5\\times 10^{-5}$, obtained by the big bang nucleosynthesis observation \\cite{bbn} and the cosmic microwave background radiation observation \\cite{cmb3}.     In this paper, we shall focus on the detection of relic gravitational waves by the cosmic microwave background (CMB) radiation observations. The RGWs leave well understood imprints on the anisotropies in temperature and polarization of CMB \\cite{cmb,cmb2}. The theoretical analysis of these imprints along with the data (including $T$, $C$, $E$, $B$) from CMB experiments allows one to determine the RGW background by constraining the parameters: the tensor-to-scalar ratio $r$ and the spectral index $n_t$. The current observations of CMB by WMAP satellite place an interesting bound $r\\le 0.20$ \\cite{wmap7} by assuming $n_t=-r/8$, which has been generalized in \\cite{zhao2010}. These bounds are equivalent to the constraints on the energy density $\\Omega_{\\rm gw}(f)$ of relic gravitational waves at the lowest frequency range $f\\sim10^{-17}$Hz.      Detecting the relic gravitational waves remains one of the most important tasks for the upcoming CMB observations (see \\cite{task} for reviews). Due to the various large contaminations, in the near future, we can only expect to detect a signal of RGWs in a relative low signal-to-noise ratio ($S/N$). This result would guide the far future detections.  As for the whole data analysis, we expect that, the maximum value of the parameters in the posterior possibility density function (pdf) is unbiased for the `true' values of the parameters, which is auto-satisfied when the $S/N$ is high. However when $S/N$ is low, the maximum values of the parameters sometimes lead to a biased guide for the `true' values, which can be generated either by some systematics or by the statistics, and should be avoided in any data analysis.  In this Letter, we will point out that,  a statistical bias could exist in the CMB data analysis for the detection of RGWs.  We also explain why the bias does exist, and suggest the way to avoid it. ",
        "conclusions": "In this Letter, we find a statistical bias in the CMB data analysis for the detection of RGWs, when the signal-to-noise ratio is not very high. This could overlook the signal of RGWs in the CMB data analysis. We explain why this bias does exist by the analytical approximation of the likelihood function, and also find this bias can be elegantly avoided by adopting the orthogonalized parameters $r$ and $z\\equiv rn_t$, instead of the general parameters $r$ and $n_t$.  In the end, we should emphasize that a similar statistical bias might exist for any data analysis \\cite{prior}, which should be carefully treated.       ~    {\\bf Acknowledgement:} The author thanks D.Baskaran, L.P. Grishchuk, P.Coles,  H.Chiaka, S.Gupta for helpful discussions. This work is supported by NSFC grants Nos. 10703005, 10775119 and 11075141."
    },
    "1107/1107.1267_arXiv.txt": {
        "abstract": "This paper presents a post-Newtonian (PN) template family of gravitational waveforms from inspiralling compact binaries with non-precessing spins, where the spin effects are described by a single ``reduced-spin'' parameter.  This template family, which reparametrizes all the spin-dependent PN terms in terms of the leading-order (1.5PN) spin-orbit coupling term \\emph{in an approximate way}, has very high overlaps (fitting factor $> 0.99$) with non-precessing binaries with arbitrary mass ratios and spins. We also show that this template family is ``effectual'' for the detection of a significant fraction of generic spinning binaries in the comparable-mass regime ($m_2/m_1 \\lesssim 10$), providing an attractive and feasible way of searching for gravitational waves (GWs) from spinning low-mass binaries. We also show that the secular (non-oscillatory) spin-dependent effects in the phase evolution (which are taken into account by the non-precessing templates) are more important than the oscillatory effects of precession in the comparable-mass ($m_1 \\simeq m_2$) regime. Hence the effectualness of non-spinning templates is particularly poor in this case, as compared to non-precessing-spin templates. For the case of binary neutron stars observable by Advanced LIGO, even moderate spins ($\\LNh \\cdot \\bS/m^2 \\simeq 0.015 - 0.1$) will cause considerable mismatches ($\\sim$ 3\\% -- 25\\%) with non-spinning templates. This is contrary to the expectation that neutron-star spins may not be relevant for GW detection. ",
        "introduction": "Coalescing compact binaries consisting of stellar-mass/intermediate-mass black holes and/or neutron stars are among the most promising sources of gravitational waves (GWs) for the interferometric GW detectors like LIGO, Virgo and GEO\\,600. Compact binary systems consisting of neutron stars or black holes can be produced in a variety of astrophysical scenarios, which can be broadly classified into two main classes: 1) Isolated binary evolution in which two massive stars constituting a binary undergo successive supernova explosions without disrupting the binary orbit~\\cite{lrr-2006-6} 2) Dynamical formation scenarios in which two compact objects form a bound orbit due to dynamical interaction in dense stellar environments~\\cite{O'Leary:2007qa}. Once formed, the binary loses orbital energy and angular momentum through GW emission and starts to inspiral and finally the binary components merge with each other.  While spin measurements of stellar-mass black holes have indicated that many black holes may have very high spins ($||\\bS||/m^2 \\sim 0.2 - 0.98$)~\\cite{McClintock:2011zq}, most of the observed neutron stars are found to be weakly spinning~\\cite{2005AJ....129.1993M}. Additionally, for neutron stars, there is a theoretical upper limit ($||\\bS||/m^2 \\sim 0.7$) on the spin rate beyond which the neutron star is gravitationally unstable~\\cite{1994ApJ...424..823C,Lattimer:2006xb}. In the case of compact binaries formed in isolated evolution, it is likely that the spins will be nearly aligned to the orbital angular momentum~\\cite{Kalogera:2004pr}. But the distribution of the spin tilt angles (angle between the spin and orbital angular momentum) depends on the distribution of supernova kicks~\\cite{1996Natur.381..584K}.  On the other hand, in the case of binaries formed in dynamical interactions, there is no prior reason to expect any spin alignments. If the spins are misaligned with the orbital angular momentum, the general relativistic spin-orbit and spin-spin coupling will cause the spins to precess around the (nearly fixed) direction of the total angular momentum~\\cite{PhysRevD.49.6274}. The complexity of the binary's dynamics will be encoded in the GW signals observed by a detector.  The GW signals from compact binaries buried in the noisy data of interferometric detectors are best extracted by employing the technique of \\emph{matched filtering}, which involves cross correlating the data with theoretically calculated templates of the expected GW signals. Theoretical GW templates can be constructed by solving Einstein's equations using analytical approximation methods and/or numerical techniques. When the compact objects are well separated and are slowly moving ($v/c \\ll 1$) gravitational waveforms can be computed using the post-Newtonian (PN) approximation to General Relativity~\\cite{Blanchet:LivRev}. The PN approximation breaks down as the compact objects reach the ultra-relativistic regime, and an accurate description of the merger process requires exact solutions of Einstein's equations which can only be obtained by large-scale numerical simulations~\\cite{Pretorius:2005gq,Campanelli:2005dd,Baker05a}.  Analytical GW templates (parameterized by the masses and spins) can be constructed by combining PN calculations with numerical-relativity simulations~(see, e.g., \\cite{Buonanno:2007pf,Damour:2009kr,Ajith:2007qp,Ajith:2007kx,Santamaria:2010yb,Sturani:2010ju}). In the case of ``low-mass'' ($m_1 + m_2 \\lesssim 12 M_\\odot$~\\cite{Ajith:2007xh,Robinson:2008un,Buonanno:2009zt}) binaries where the signal-to-noise ratio (SNR) is almost entirely contributed by the inspiral portion of the signal, it is sufficient to model only the inspiral accurately, where the PN approximation holds. Although the expected form of the signals can be computed as a function of the source parameters, the parameters of a signal that is buried in the noise is not known \\emph{a priori}. Thus the data is cross correlated with a ``bank'' of theoretical templates corresponding to different (astrophysically plausible) values of physical parameters. A geometrical formalism has been developed for ``laying down'' templates in the parameter space of compact binaries~\\cite{Sathyaprakash:1991mt,Owen:1995tm}. Parameters of the templates are chosen in such a way that the loss of SNR due the mismatch between two neighboring templates is less than an acceptable value, while at the same time keeping the total number of templates employed in the search computationally tractable.  However, templates for spinning binaries are characterized by a large number of parameters (two for the masses, six for the spins, two for the inclination and polarization angles), complicating placement of templates in the bank and considerably increasing the computational cost of searches. Furthermore, many different spin configurations are known to be degenerate, making it unnecessary to employ templates corresponding to \\emph{all} parameters.  The idea of constructing an \\emph{effective} template family parameterized by a smaller number of parameters that has high enough overlaps with the expected signals was first proposed by Apostolatos~\\cite{PhysRevD.54.2438}, who introduced a modulational sinusoidal term in the frequency-domain phase of the templates to capture the oscillatory effects of precession. However, it was found that this template family fails to capture the target signals with sufficient efficiency~\\cite{PhysRevD.67.042003}. Later, Buonanno, Chen and Vallisneri (BCV)~\\cite{BCV2} proposed a phenomenological template family that has better overlaps with the target signals, by introducing modulation effects in both the frequency-domain- amplitude and phase of the templates. Several of these phenomenologically introduced parameters could be searched over analytically, thus reducing computational costs significantly. A search for GWs from spinning binaries was performed using the data from LIGO's third science run, employing BCV templates~\\cite{PhysRevD.78.042002,Jones:2010iv}. Interestingly, the sensitivity of this search towards spinning binaries was found to be not any better than a search employing non-spinning templates. Indeed, BCV have cautioned that the increased degrees of freedom associated with the detection statistic (due to the analytical maximization of several phenomenological parameters) will also increase the false alarm rate. Later Van Den Broeck~\\etal~\\cite{VanDenBroeck:2009gd} demonstrated that, while BCV templates have high overlaps with the target signals, the increased degrees of freedom associated with the detection statistic and the lack of good methods for vetoing non-Gaussian detector glitches that mimic the expected signals (``signal-based vetoes'') resulted in producing a much higher false alarm rate in the presence of non-Gaussian noise, and hence negated the advantages. They argued that the standard non-spinning frequency-domain 3.5PN templates have higher detection efficiency (at a given false alarm rate) towards spinning binaries than the current implementation of the phenomenological BCV template bank.  BCV~\\cite{BCV2} also proposed a physical template family, which presumably will not suffer from this limitation. This template family, which assumes that only one compact object has significant spin, is parametrized by four parameters that need to be searched over (two masses, magnitude of the single spin, spin tilt angle). This template family was extensively studied by Pan~\\etal~\\cite{Pan:2003qt} and Buonanno~\\etal~\\cite{Buonanno:2004yd}, where they demonstrated that the template family is \\emph{effectual}~\\cite{DIS98} in detecting single-spin as well as double-spin binaries using Initial detectors. They also demonstrated the possibility of reducing the number of spin parameters to one in certain regions in the parameter space. This template family has been implemented in the LIGO-Virgo search pipelines~\\cite{FaziThesis,Harry:2010fr,Harry:2011qh}. Another enhancement of the non-spinning frequency-domain template family which extends the symmetric mass ratio ($\\eta = m_1 m_2/m^2$) to unphysical values ($\\eta > 0.25$) has also been proposed as a detection template family for spinning binaries~\\cite{Boyle:2009dg,Aylott:2009ya,Lundgren}.  In this paper we propose a frequency-domain PN template family characterized by a \\emph{single} ``reduced-spin'' parameter (and two masses) for the case of binaries with non-precessing spins (spins aligned/anti-aligned with the orbital angular momentum). This template family, which reparametrizes all the spin-dependent PN terms in terms of the leading-order (1.5PN) spin-orbit coupling term \\emph{in an approximate way}, has very high overlaps with non-precessing binaries with arbitrary mass ratios and spins. We also show that this ``reduced-spin'' template family is able to capture a significant fraction of generic precessing binaries in the comparable-mass regime ($q \\equiv m_2/m_1 \\lesssim 10$). This might provide an efficient and feasible way of searching for spinning low-mass binaries in the comparable-mass regime.  This result also shows that the secular (non-oscillatory) spin-dependent effects in the phase evolution (which are taken into account by the non-precessing templates) are more important than the oscillatory effects in the case of (nearly) equal-mass binaries even with generic spins.  This paper is organized as follows: Section~\\ref{sec:summary} summarizes the main findings of this paper as well as lists the limitations of this work. Section~\\ref{sec:targetWaves} provides a description of the PN waveforms from precessing binaries computed in the adiabatic approximation, which are assumed to be the ``target signals'' that we want to detect. Section~\\ref{sec:Detection} is a brief introduction to the GW data analysis of inspiralling compact binaries. In Section~\\ref{sec:EffSpinTempl}, we construct the frequency-domain PN template family parametrized by a reduced spin parameter and demonstrate the \\emph{effectualness}~\\cite{DIS98} of the template family in detecting binaries with non-precessing spins. In Section~\\ref{sec:Precession} we investigate the effectualness of the template family in detecting precessing binaries. We use geometrical units throughout the rest of the paper: $G = c = 1$. ",
        "conclusions": "In this paper, we presented a PN detection template family of gravitational waveforms from inspiralling compact binaries with non-precessing spins. The waveforms are reparametrized in such a way that all spin-dependent terms are described by a single ``reduced-spin'' parameter in an approximate fashion. We have shown that the template family has very high overlaps with non-precessing binary signals with arbitrary spins and mass ratios. The family is also effectual for the detection of a significant fraction of precessing binaries in the comparable-mass regime. This is due to the fact that, in the comparable-mass regime, the secular spin-dependent effects dominate the phase evolution (not the oscillatory effects), which are captured by the non-precessing templates. Since non-spinning templates neglect this effect, a significant fraction of SNR can be lost if non-spinning templates are used in the search for spinning binaries.  The simple, closed-form expression of this frequency domain template family makes it very easy to implement this in search pipelines, and will provide considerable improvement over the non-spinning templates. Since a parameter-space metric~\\cite{Owen:1995tm} can be easily worked out for this waveform family, the current techniques of template placement and multi-detector coincidence tests~\\cite{Robinson:2008un} can be readily used in such a search. The parameterization of the inspiral waveforms using the reduced spin parameter $\\chi$ can also be extended to the construction of inspiral-merger-ringdown waveform templates for binary black holes, such as the ones presented in~\\cite{Ajith:2009bn,Santamaria:2010yb}.  \\smallskip"
    },
    "1107/1107.3748_arXiv.txt": {
        "abstract": "\\grb{} is among the most energetic bursts detected and the most extensively observed so far. Nevertheless, unresolved issues are still open in the literature on the physics of the afterglow and on the GRB environment. In particular, \\grb{} was claimed to have exploded in a galactic halo environment, based on the uniqueness of the optical spectrum and the non-detection of an underlying host galaxy. In this work we collect all publicly available data and address these issues by modelling the NIR-to-X-ray spectral energy distribution (SED) and studying the high signal-to-noise VLT/FORS afterglow spectrum in comparison with a larger sample of GRB absorbers. The SED reveals a synchrotron cooling break in the UV, low equivalent hydrogen column density and little reddening caused by a LMC- or SMC-type extinction curve. From the weak \\mgii{} absorption at $z=1.5477$ in the spectrum, we derived $\\log N$(\\mgii{}$)=12.96^{+0.13}_{-0.18}$ and upper limits on the ionic column density of several metals. These suggest that the GRB absorber is most likely a Lyman limit system with a $0.03\\,Z_\\odot <Z<1.3\\,Z_\\odot$ metallicity. The comparison with other GRB absorbers places \\grb{} at the low end of the absorption line equivalent width distribution, confirming that weak spectral features and spectral-line poor absorbers are not so uncommon in afterglow spectra. Moreover, we show that the effect of photo-ionization on the gas surrounding the GRB, combined with a low $N$(\\hi{}) along a short segment of the line of sight within the host galaxy, can explain the lack of spectral features in \\grb{}. Finally, the non-detection of an underlying galaxy is consistent with a faint GRB host galaxy, well within the GRB host brightness distribution. Thus, the possibility that GRB 070125 is simply located in the outskirts of a gas-rich, massive star-forming region inside its small and faint host galaxy seems more likely than a gas-poor, halo environment origin. ",
        "introduction": "The progenitors of long gamma ray bursts (GRBs) are believed to be rapidly rotating massive stars \\citep[for a review, see][]{Woosley93,MacFadyen99}. The observations of some SNe associated with GRBs have corroborated this hypothesis \\citep[][and references therein]{Woosley06,Hjorth11}, implying that long GRBs are located in star-forming regions. Within the fireball model \\citep[see][for a review]{Piran04,Meszaros06} the GRB afterglow is produced by the interaction of an expanding blast wave with the surrounding medium. The external shocks are responsible for the synchrotron radiation, from X-rays to optical and down to radio frequencies as the blast wave decelerates. The spectral energy distribution (SED) and its evolution is described by the fireball model \\citep{Sari98,Granot02}. Modelling the SED can simultaneously provide information on the intrinsic spectral shape of the GRB, i.e. the fireball physical parameters, and on the GRB local and host galaxy environment. In the latter case, this includes the circumburst density or the line of sight dust extinction, depending on whether the broadband blast wave, down to radio frequencies, is modelled \\citep[e.g.,][]{Panaitescu02,Bjornsson04,Postigo05,Vanderhorst05,Johannesson06,Rol07} or the optical-to-X-ray SED \\citep{Starling07,Starling08,Schady10, Zafar11}. The interstellar medium of the GRB host galaxy can also be investigated through absorption line spectroscopy of the afterglow, providing the chemical composition of different ions in the absorbing gas and possibly its metallicity \\citep[e.g.,][]{Vreeswijk04,Fynbo06,Savaglio06,Prochaska07,Ledoux09}.  \\grb{} is among the most energetic bursts detected so far, both in the prompt high energy release \\citep[with an isotropic emitted energy $E_{\\rm iso}\\sim10^{54}$ erg:][]{Bellm08} and its afterglow \\citep[with a collimation-corrected blast wave kinetic energy of more than $10^{52}$ erg:][]{Chandra08}. The optical afterglow is amongst the brightest at around one day after the burst, when redshift-corrected \\citep[$R_c\\sim17$ shifted at $z=1$:][]{Updike08,Kann10}. The brightness of the event allowed a global monitoring campaign of the afterglow resulting in one of the largest datasets collected so far, covering radio, mm, NIR, optical, UV and X-rays \\citep{Cenko08,Chandra08,Dai08,Updike08,Kann10}. Despite the observational effort, there is still much debate over both the jet properties and the type of environment in which \\grb{} was situated.  Based on the absence of strong spectral features and of a bright underlying host, \\grb{} was reported to be the first long GRB to have exploded in a galactic halo environment. This would pose a challenge to modelling of progenitor objects. In particular, the GRB site was suggested to be in a compact stellar cluster that resulting from the interaction of two blue faint galaxies at a projected distance of $\\approx 27$ and $\\approx 46$ kpc\\footnote{Derived from the $3.2\\arcsec$ and $5.5\\arcsec$ angular distance between the afterglow and the two putative host galaxies.} \\citep{Cenko08}. However, a very limited sample of GRB afterglows was used as reference to argue the peculiarity of the spectrum (few spectral features, weak \\mgii{} absorption and low inferred $N$(\\hi{})). By contrast, from a partial broadband SED analysis, \\citet{Chandra08} found a very high circumburst density ($n\\approx50$ cm$^{-3}$, from the kinetic energy and $n\\approx40$ cm$^{-3}$, from the broadband modelling of the SED). However, in their data it was not possible to distinguish between a constant or a wind-like density profile nor to firmly constrain the putative (chromatic) jet break and the inverse Compton influence on it. On the other hand, \\citet{Dai08} claimed a ``textbook'' achromatic jet break at 5.8 days after the burst. Another open issue is the presence or absence of a synchrotron cooling break in the SED between the optical and X-rays and the agreement or not of the data with the fireball closure, as discussed in \\citep{Updike08}.   In this work, we collect the multitude of observations spread through the literature and analyse them together in order to obtain a coherent picture of the GRB environment. In particular, we investigate the location of a synchrotron cooling break and the environment dust and metal content by studying the NIR-to-X-ray SED at different epochs. Furthermore, we analyse the very high signal-to-noise (S/N) VLT/FORS spectrum of the afterglow and discuss the possible implications for the absorber environment. We then compare the spectral features to the large sample of GRB afterglows with redshift collected in \\citet{Fynbo09} and two sub-samples of GRBs potentially similar to \\grb{}. The observations and data analysis are presented in Sec. \\ref{sec obs} and Sec. \\ref{sec analysis} respectively, the SED modelling in Sec. \\ref{sec mod} and the discussion in Sec. \\ref{sec discussion}. Our conclusions are summarized in Sec. \\ref{sec conclusions}.  Throughout the paper we use the convention $F_\\nu \\left(t\\right) \\propto t^{-\\alpha}\\nu^{-\\beta}$ for the flux density  of the afterglow, where $\\alpha$ is the temporal slope and $\\beta$ is the spectral slope. We refer to the Solar abundances measured by \\citet{Asplund09} and adopt cm$^{-2}$, as the linear unit of column densities, $N$. Hereafter, we assume a standard $\\Lambda$CDM cosmology with $H_0$ = 70.4 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.27 and $\\Omega_{\\Lambda}$ = 0.73 \\citep{Jarosik11}. ",
        "conclusions": "\\label{sec conclusions}  We collected all observational data of \\grb{} available in the literature, in order to investigate the environmental properties of the burst, and in particular to test the putative halo origin of this burst based on the few weak absorption lines in the spectrum and the absence of an underlying bright host \\citep{Cenko08}.  The NIR-to-X-ray SED at 1.9, 2.5 and 2.9 days after the burst revealed a synchrotron cooling break in the UV. Along the line of sight in the host galaxy, little reddening ($E(B-V) = 0.016\\pm0.007$) is caused by a LMC- or SMC-type dust, while a low equivalent hydrogen column density cannot be constrained better than $N$(H)$\t=\t15.5^{+13.9}_{-14.9} \\times 10^{20}$ cm$^{-2}$. The NIR-to-X-ray SED results are inconsistent with broadband SED analyses including radio observations, indicating that the assumptions in modelling efforts for \\grb{} so far have probably been too simplified. More complex afterglow models should be explored, in particular to unify the light curves across the broadband spectrum with the NIR-to-X-ray SED.  The analysis of the high S/N, but rather featureless FORS spectrum showed weak \\mgii{} and \\civ{} in absorption, providing $\\log N$(\\mgii{}$)=12.96^{+0.13}_{-0.18}$. The constraint on the \\mgii{} column density, together with the upper limits for the other ions and the evidence for low metallicity from the SED, suggested that the GRB absorber is most likely a $0.03\\,Z_\\odot <Z<1.3\\,Z_\\odot$ Lyman limit system, also observed along the line of sight to five other GRBs. The comparison of the \\grb{} measurements with the rest-frame $EW$s of the GRB afterglow sample of \\citet{Fynbo09} revealed that \\grb{} is not unique, but, simply falls at the low end of the $EW$ distribution. Furthermore, we selected two sub-samples of afterglows possibly similar to \\grb{}, the spectral-line poor and the low $N$(H). Most spectral-line poor afterglows show a low $A_V$ and low $N$(H), suggesting that a short distance in the host galaxy is traversed by the line of sight. In any case, a quite featureless spectrum is not as uncommon for GRB afterglows as previously claimed.  We demonstrated that the few and weak features in the spectrum of \\grb{} can be the result of the particularly intense afterglow radiation ionizing a low $\\log N$(\\hi{})~$\\sim19$ through a line of sight traversing not more than about 0.5 kpc of its faint host galaxy. We showed that few and weak spectral features are not so uncommon in GRB afterglows, including the nearby GRB\\,101219B at $z=0.55$, with associated SN, for which a sub-luminous host galaxy has been claimed. Finally, the non-detection of a host galaxy at the \\grb{} position indicates that an underlying galaxy, if present, must be among the faintest GRB host galaxies.      Given that about half of the GRBs in the spectral-line poor sub-sample have a detected host galaxy, it seems likely that these bursts are also located in the outskirts of a gas-rich, massive star-forming region inside its small and faint host galaxies, rather than in a halo environment."
    },
    "1107/1107.3781_arXiv.txt": {
        "abstract": "{The Galactic Center offers unique opportunities to study stellar and bow-shock polarization effects in a dusty environment.} {The goals of this work are to provide near-infrared (NIR) polarimetry of the stellar sources in the central parsec at the resolution of an 8m telescope for the first time, along with new insights into the nature of the known bright bow-shock sources.} {We use adaptive-optics assisted observations obtained at the ESO VLT in the H- and Ks-band, applying both high-precision photometric methods specifically developed for crowded fields and a newly established polarimetric calibration for NACO to produce polarization maps of the central 3''$\\times$19'', in addition to spatially resolved polarimetry and a flux variability analysis on the extended sources in this region.}{We find foreground polarization mainly parallel to the Galactic plane, with average values of (4.6 $\\pm$ 0.6)\\% at $26^{\\circ} \\pm 6^{\\circ}$ (Ks-band) and (9.3 $\\pm$ 1.3)\\% at $20^{\\circ} \\pm 6^{\\circ}$ (H-band) in the center of the field-of-view (FOV). Further away from the center, we find higher polarization degrees and steeper polarization angles: (7.5 $\\pm$ 1.0)\\% at $11^{\\circ} \\pm 6^{\\circ}$ (Ks-band) and (12.1 $\\pm$ 2.1)\\% at $13^{\\circ} \\pm 6^{\\circ}$ (H-band). $p_H / p_{Ks}$ peaks at $1.9 \\pm 0.4$, corresponding to a power law index for the wavelength dependency of $\\alpha = 2.4 \\pm 0.7$. These values also vary over the FOV, with higher values in the center. This is indicative of the influence of local effects on the total polarization, possibly dichroic extinction by Northern Arm dust. The two extended sources IRS~21 and 1W show similar intrinsic polarization degrees of 6.1 resp. 7.8\\% (Ks) and 6.9 (H, only 1W) at polarization angles coincident with previous NIR and mid-infrared (MIR) findings, both in total and spatially resolved. The spatial polarization pattern of both sources points to scattering on aligned elongated dust grains as the major source of intrinsic polarization, and matches the known orientation of the magnetic field. Our data also allow us to separate the bow shock of IRS~21 from the central source for the first time in the Ks-band, finding the apex north of the central source and determining a standoff distance of $\\sim$400 AU, which matches previous estimates. This source also shows a $\\sim50$\\% increase in flux in the NIR over several years.}{} ",
        "introduction": "The center of the Milky Way is located at a distance of $\\sim8.0$ kpc \\citep{ghez2008,gillessen2009}. It contains a nuclear stellar cluster (NSC) with a $\\sim4.0 \\times 10^6$ M$_{\\sun}$ super-massive black hole at its dynamical center \\citep{eckart2002,schoedel2002,schoedel2003,ghez2003,ghez2008,gillessen2009}. This cluster shows similar properties as the NSCs found at the dynamical and photometric centers of other galaxies \\citep{boeker2010,schoedel2010c}.\\\\ Less than a decade after the first near-infrared (NIR) imaging observations of the Galactic center (GC) by \\cite{becklin_neugebauer1968}, the polarization of 10 sources within $\\leq$2 pc of Sagittarius A* was measured by \\cite{capps1976,knacke1977}, observing in the K-, L- and 11.5 $\\mu$m-band. These observations revealed similar polarization degrees and angles for the four sources observed in the K-band, with the polarization angles roughly parallel to the galactic plane \\citep[$\\sim$4\\% at 15$^{\\circ}$ East-of-North, while the projection of the galactic plane is at a position angle of $\\sim$31.4$^{\\circ}$ at the location of the GC, see ][]{reid2004}. This was interpreted as polarization induced by aligned dust grains in the Milky Way spiral arms along the line-of-sight (LOS). The 11.5 $\\mu$m polarization, however, was found to be almost perpendicular to the galactic plane and was therefore classified as intrinsic. The values found for the L-band showed intermediate values, which was attributed to a superposition of both effects. The GC therefore offers the possibility of studying both interstellar polarization and the properties of intrinsically polarized sources.\\\\ \\cite{kobayashi1980} conducted a survey of K-band polarization in a much wider field of view (7'$\\times$7'), finding largely uniform polarization along the galactic plane. \\cite{lebofsky1982} confirmed these findings in the H- and K-band for 17 sources in the central cluster. The latest H-band survey was conducted by \\cite{bailey1984}, who examined $\\sim$10 sources with a 4.5'' resolution, finding similar results. While these surveys could only resolve a small number of sources in the central region, higher resolution observations (0.25'') enabled \\cite{eckart1995} to measure the polarization of 160 sources in the central 13''$\\times$13'', while \\cite{ott1999} examined $\\sim$ 40 bright sources in the central 20''$\\times$20'' at 0.5'' resolution. These two studies confirmed the known largely uniform foreground polarization, but already revealed a more complex picture: individual sources showed different polarization parameters, such as a significantly higher polarization degree (IRS~21).\\\\ These results already illustrated that in addition to the foreground polarization, intrinsic polarization plays a role in the Ks-band. Especially sources embedded in the Northern Arm and other bow-shock sources (see Fig.\\ref{FigLband} for an overview of the bright bow-shock sources in the central 20'') show signs of intrinsic polarization. Among these objects, IRS~21 shows the strongest total Ks-band polarization of a bright GC source detected to date \\citep[$\\sim$ 10-16\\% at 16$^{\\circ}$]{eckart1995,ott1999}. \\begin{figure*}[!t] \\centering \\includegraphics[width=\\textwidth,angle=0, scale=1.0]{ppRMB2011f1.eps} \\caption{\\small Ks-band field-of-view, one channel of the Wollaston prism.} \\label{FigFOVwoll} \\end{figure*} \\cite{tanner2002} described IRS~21 as a bow-shock most likely created by a mass-losing Wolf-Rayet star. The observed polarization is a superposition of foreground polarization and source intrinsic polarization. It is still unclear what process(es) cause the latter component: (Mie)-scattering in the dusty environment of the Northern Arm and/or emission from magnetically aligned dust. The latter is known to occur at 12.5 $\\mu$m \\citep{aitken1998}, but should be negligible at shorter wavelengths.\\\\ \\begin{figure}[!b] \\centering \\includegraphics[width=\\textwidth,angle=0, scale=0.48]{ppRMB2011f2.eps} \\caption{\\small Lp-band image of the innermost 20'' of the GC (ESO VLT NACO image, program 179.B-0261(A)). Prototypical (bow-shock) sources and location of the Northern Arm are indicated. Stellar sources are marked with stars, bow-shocks are indicated by triangles.} \\label{FigLband} \\end{figure} In general, three effects can produce polarized NIR radiation in GC sources: (re)emission by heated, non-spherical grains, scattering (on spherical and/or aligned non-spherical grains) and dichroic extinction by aligned dust grains. The first two cases can be regarded as intrinsic to the source for our purposes, thereby allowing conclusions about the source itself and its immediate environment, while the third effect is the result of grain alignment averaged along the LOS. For sources enclosed in an optically thick dust shell, dichroic extinction can also contribute significantly as a local effect \\citep[see e.g.][]{whitney2002}.\\\\ As the basic mechanism that could cause the observed large-scale grain alignment, \\cite{davis_greenstein1951} suggested paramagnetic dissipation, which basically aligns the angular momentum of spinning grains with the magnetic field. But even almost 60 years later, the problem of grain alignment is by no means completely solved, and it remains difficult to reach exact conclusions for dust parameters and magnetic field strength, but at least determining the magnetic field orientation is possible. If the parameters change along the LOS, this further complicates the issue. See e.g. \\cite{purcell1971}, \\cite{lazarian2003}, cite{lazarian2007} and references therein for a review of the different possible causes of grain alignment expected to be relevant in different environments.\\\\ This makes it possible to use polarimetric measurements to map at least the direction of the magnetic fields responsible for dust alignment through the Davis-Greenstein effect, as e.g. \\cite{nishiyama2009,nishiyama2010} showed for the innermost 20' resp. 2$^{\\circ}$ of  the GC, but these studies did not cover the central parsec owing to insufficient resolution.\\\\ Observations of the Galactic Center suffer from strong extinction caused by dust grains on the LOS, with values of up to $A_V=40$ mag at optical wavelengths (or even higher values of up to 50 mag if a steeper extinction law is assumed) and still around 3 mag in the Ks-band \\citep[e.g.][]{scoville2003,schoedel2010b}. The aligned interstellar dust grains responsible for the polarization cause extinction as well, but non-aligned grains can also contribute. Therefore, the same particles are not necessarily responsible for both effects \\citep[e.g.][]{martin1990}. Universal power laws have been claimed for NIR extinction by \\cite{draine1989} and polarization \\citep{martin1990}, who also showed that the law applicable to polarization in the optical domain \\citep{serkowski1975} is a poor approximation in the NIR. The power law indices presented in these works for the extinction and the polarization power law are almost the same (1.5-2.0). It also appears that there is a correlation between the measured extinction of an intrinsically unpolarized source and foreground polarization \\citep{serkowski1975}, but this relation is quite complex. In the light of new results for the extinction law, which seem to deviate consistently from the Draine law \\citep[e.g.][]{gosling2009,schoedel2010b}, new polarimetric measurements may be useful to further clarify the relation between extinction and polarization. \\begin{figure}[!t] \\centering \\includegraphics[width=\\textwidth,angle=0, scale=0.48]{ppRMB2011f3.eps} \\caption{\\small Total relative flux errors of the sources detected in the Ks- (upper frame) resp. H-band (lower frame). Only the 0$^{\\circ}$ channel is shown here, but the error distribution is very similar in the other three channels. The values are based on variations between the fluxes measured for the same source at different dither positions. The red lines denote the 16/18 mag resp. 3/6\\% upper limits (see Appendix \\ref{SectionErrors}).} \\label{FigPhoterrsHK} \\end{figure} The central parsec of the GC is a well-suited but challenging environment to study this relation, because it contains a huge number of sources that exhibit large variances in extinction \\citep[1-2 mag, see][]{buchholz2009,schoedel2010b}, which is produced along a long line-of-sight by a great number of dust clouds with possibly different grain alignment and composition. We are using a new Ks-band extinction map of the central parsec recently been presented by \\cite{schoedel2010b}.\\\\ The aims of this work are to present the first map of stellar H- and Ks-band polarization of the central few arcseconds, a highly crowded region with more than 10 sources per arcsec$^2$. In order to avoid crowding and to resolve the sources individually, the resolution of an 8m class telescope is required. Especially in the H-band, no polarization measurements exist so far for individual resolved sources in the central parsec. In addition, we will present the first spatially resolved polarimetric measurements on IRS~1W and IRS~21, also resolving the bow-shock structure of IRS~21 for the first time in the Ks-band. Furthermore, we will conduct a variability analysis on the extended sources IRS~1W, 5, 10 and 21 in the H-, Ks- and L-band.\\\\ In \\S \\ref{SectObs} we discuss the photometric methods used and the calibration applied to the data. Our results will be presented in \\S \\ref{SectResults}, followed by a summary and discussion of their implications in \\S \\ref{SectDiscussion}. ",
        "conclusions": "\\\\ \\begin{enumerate} \\item Like several previous polarimetric surveys of the central parsec at much lower resolution, we find that the polarization follows the Galactic plane in general. This confirms that the polarization can indeed be attributed to aligned dust grains between the Galactic Center and the observer (foreground polarization). We find systematic and individual deviations from this general pattern, however. Over our FOV we find that the polarization degrees and angles change toward the eastern region around IRS~1. This behavior occurs in the H- and the Ks-band data, with stronger relative deviations in the Ks-band.  That can be expected if this is indeed a local effect produced by dust grains in the Northern Arm, because typical temperatures and grain sizes in this region lead to a much more significant impact of radiation from this grain population at longer wavelengths. This is supported by a systematic change of $\\frac{p_H}{p_{Ks}}$ over the FOV, with values of 2.0$\\pm$0.3 in the center and 1.6$\\pm$0.3 in the Northern Arm region. The higher of these values leads to a polarization maximum at $\\sim 0.7 \\mu$m, which matches previous findings of $\\lambda_{max} \\sim 0.8 \\mu$m by \\cite{bailey1984}, based on JHK observations of a more extended region in the GC. This value represents dust parameters, specifically typical grain sizes, and it can be expected that local contributions by distinct dust masses lead to a different value than the average along the LOS. Unfortunately, we cannot derive any more constraints for the local dust based only on these two bands, other than that the dust grains responsible for the local effects must be elongated and aligned \\citep[see][]{aitken1998}, and that their temperature and size must be sufficient for their influence in the H/Ks-band. \\item We find a correlation between the spatially variable extinction toward the central parsec and the polarization efficiency, $\\frac{p_{\\lambda}}{A_{\\lambda}}$. This relation appears to be different for the two groups of sources we find in our FOV (sources only affected by LOS polarization respectively sources where local contributions are significant): for the foreground-polarized sources, we find a power law relation with a similar power law index for the H- and the Ks-band ($\\beta_H = -0.6 \\pm 0.6$, $\\beta_{Ks} = -0.4 \\pm 0.4$ ). This is compatible with the measurements resp. the model by \\cite{jones1989}, who assume a combination of a constant and a random component of the galactic magnetic field along the LOS to align the dust grains. For the population of higher polarized sources, we find a higher polarization efficiency in general, while the power law index for both bands resembles that of the sources with lower polarization (with $\\beta_H = -0.4 \\pm 0.7$ and $\\beta_{Ks} = -0.5 \\pm 1.5$, although it has to be cautioned that these fits are poor because of the small number of sources). The higher polarization efficiency points to the influence of local dust grains, which only contribute a small amount to the total extinction. But because these grains are aligned, they add a significant component to the total polarization, because this alignment is not (partially) compensated by averaging over other dust populations with different alignment as is the case for every individual dust configuration along the LOS. These findings further support our previous conclusion that the influence of local dust is quite significant not only in the MIR, but also in the H- and Ks-band. \\item Intrinsic polarization takes place in several sources in the central parsec, not only at longer wavelengths as shown first by \\cite{knacke1977}, but also in the H- and Ks-band \\citep[e.g.][]{ott1999}. With the M\\\"uller calculus we can isolate the intrinsic component for point sources and maps of extended features. The resulting total intrinsic polarization angles for IRS~21 and IRS~1W agree very well with what can be expected from source morphology and relative motion of gas/dust in the Northern Arm and the sources themselves. We find very similar intrinsic polarization degrees for both sources, with (7.8 $\\pm$ 0.5) \\% at (-75 $\\pm$ 5)$^{\\circ}$ for IRS~1W and (6.1 $\\pm$ 0.5) \\% at (5 $\\pm$ 5) for IRS~21 (both Ks-band). Contrary to the wavelength dependency seen in the foreground polarization, we find a slightly lower H-band polarization degree for IRS~1W compared to the Ks-band: (6.9 $\\pm$ 0.5) \\% at (-73 $\\pm$ 5)$^{\\circ}$. That the central source contributes a larger amount of flux in the H-band compared to the extended component, while the intrinsic polarization mostly stems from the bow-shock, explains this discrepancy.\\\\ The spatially resolved intrinsic polarization patterns of both bright extended sources in our FOV do not show a centrosymmetric shape as would be expected for Mie scattering on spherical grains. Instead, the pattern found for IRS~1W suggests that scattering on aligned elongated grains must play a significant role, possibly combined with thermal dust emission (in the Ks-band), because the grains in the bow-shocks may reach temperatures on the order of 1000 K \\cite[see][]{moultaka2004}. Both effects would produce polarization parallel to the long axis of the grains and thus perpendicular to their angular momentum. It therefore appears that the angular momentum of the grains in the bow-shock surrounding IRS~1W is aligned parallel to the relative motion of the source through the Northern Arm material. Taking the magnetic fields in the Northern Arm into account, this is most likely caused by the field lines being warped around the flanks of the bow-shock. This in turn compresses the field lines and thus increases the local field strength, which is then sufficient for fast magnetic alignment of the grains.\\\\ For IRS~21 we find a different polarization pattern. It appears that the polarization vectors are mostly parallel to the relative motion of the source through the surrounding material. Combined with the local field orientation, this again leads to the same dominant mechanism of scattering/emission as for IRS~1W, if we assume that the field lines are compressed in front of the shock that is moving perpendicular to them. \\item After deconvolution, IRS21 clearly does not show a circular shape in the Ks-band. Instead, its shape is consistent with that of a poorly resolved bow-shock with its apex north of the central source. This is consistent with the L-band shape presented by \\cite{tanner2004} and the expectations from the proper motions of the source in relation to the streaming motion of the Northern Arm. We find a standoff distance of $\\sim$400 AU, this agrees with the upper limit given by \\cite{tanner2002}. It is difficult to constrain the nature of the central source by this parameter alone, and most stellar types with strong winds, such as Wolf-Rayet stars, would be likely candidates \\citep[see][]{tanner2002}. \\item We find a flux variability for several bow-shock sources in the central parsec: IRS~21 shows a flux increase of $\\sim$50\\% over 6 years in the H- and Ks-band (less, but still significant in the Lp-band). Especially the H-band measurement indicates that this variability must be intrinsic to the central source and not only to the envelope. The candidates suggested for the central source by \\cite{tanner2002} could all show this variability, especially a source that is currently in a mass-losing phase.\\\\ While we also find a possible flux variability for IRS~5, IRS~1W, and IRS~10W in all three bands, no clear periodicity or trend is apparent for these sources at our time resolution. We therefore cannot determine if these sources show a true intrinsic variability or not. We confirm the known periodicity of IRS~10E ($\\sim$2 yrs) in the Ks-band, finding the same variability in the H- and Lp-band as well. \\end{enumerate} Several effects appear to contribute to the observed polarization in the central parsec of the GC, in addition to the foreground polarization. The findings in our relatively narrow field-of-view alone indicate possible additional areas of interest: do the large-scale trends we find continue in other sections of the Northern Arm (and the other elements of the minispiral)? Do the other bow-shock sources, such as IRS~10W, 5, and 8 show similar polarization patterns, and what does that tell us about the local processes in the bow-shocks? Are other intrinsically polarized sources present in the field, and what type of objects would show this behavior? More observations with wider FOVs and in different bands (especially L-band, to directly examine extended dust emission e.g. in the Northern Arm) would be desirable, and the recently achieved direct polarimetric calibration of NACO offers the possibility to map more extended areas of the GC in polarization at high resolution."
    },
    "1107/1107.4959_arXiv.txt": {
        "abstract": "We combine DensePak integral field unit and TAURUS Fabry-Perot observations of 13 nuclear rings to show an interconnection between the kinematic properties of the rings and their resonant origin. The nuclear rings have regular and symmetric kinematics, and lack strong non-circular motions. This symmetry, coupled with a direct relationship between the position angles and ellipticities of the rings and those of their host galaxies, indicate the rings are in the same plane as the disc and are circular. From the rotation curves derived, we have estimated the compactness ($v^2/r$) up to the turnover radius, which is where the nuclear rings reside. We find that there is evidence of a correlation between compactness and ring width and size. Radially wide rings are less compact, and thus have lower mass concentration. The compactness increases as the ring width decreases. We also find that the nuclear ring size is dependent on the bar strength, with weaker bars allowing rings of any size to form. ",
        "introduction": "Nuclear rings in the central kiloparsecs of galaxies are usually observed as conglomerations of young massive stars in distinct compact groupings \\citep{BC96, JK05, C10}. The star forming rings are believed to form as a result of radial gas inflow towards the central region, which stagnates near dynamical resonances. In barred galaxies the inflow originates in the gravitational torques set up by the bar (e.g., \\cite{A92,  HS94, HS96, P95, JK95, BC96, R97, B02}). Some 20\\% of nuclear rings occur in non-barred galaxies \\citep{JK05, C10} and in those cases the inflow and resonances may well be related to the presence of weak ovals,  past interactions or mergers, or even strong spiral arms. A resonant origin of all nuclear rings is most plausible.  The radial location of a nuclear ring can be described by the position of the inner Lindblad resonance(s) (ILRs) \\citep{S80, Sh89, BC93, FB93, HS94, JK95, BC96}, which is typically at $\\sim$1 kpc from the galaxy center although smaller and larger nuclear rings exist \\citep{BC93, C10}. In a non-axisymmetric potential, the nuclear ring forms near the ILR(s), where the inflowing molecular gas slows down, accumulates, and initiates star formation. For barred spiral galaxies, the ring can be formed inside the large scale galactic spiral structure in the vicinity of the inner Lindblad resonances (ILRs), which occur at those radii where the bar pattern speed $\\Omega_{b}$ equals ($\\Omega - \\kappa/2$), with $\\Omega$ the angular velocity and kappa the radial epicyclic frequency. This approximation is not valid for strong bars where a non-linear analysis must be performed, but in any case serves as a useful indication of how the location of the ring is related to the underlying galaxy dynamics \\citep{JK95}.  Only a few in-depth observational studies of the kinematics of nuclear rings at high angular resolution exist, but these have revealed some consistent kinematic characteristics of nuclear rings. \\cite{B96} showed that non-circular motions near the intersection of the ring and bar dust lanes correspond to radially inflowing molecular gas that is apparently feeding the ring in NGC~4314. \\cite{RD98} concluded for NGC~1530 the existence of weak velocity transitions at the ends of the bar which are inversely correlated with strong star formation in the same areas, including the nuclear ring location. \\cite{R97} further note an increase in the residual velocity along an axis perpendicular to the bar. \\cite{JK00} (see also \\cite{JK95}) and \\cite{Z04} used CO and \\halpha\\ emission in the central region of NGC~4321 and NGC~1530, respectively, to reveal gas streaming inward towards the ring along the bar dust lanes. \\cite{J02} studied the fueling mechanisms of the ring in NGC~5248 to deduce that large star forming clusters in the ring have been triggered by a bar-driven spiral density wave. More recently, \\cite{A06} follow up on the ring in NGC~4321 to show relatively low \\hbeta\\ gas dispersion within the ring, which can be an indicator of active star formation and cool gas inflow from which the massive stars inherited the low velocity dispersion.  The existing kinematic studies provide insight into a small number of nuclear rings, but fail to reach more general conclusions applicable to the overall population. Apart from the differing characteristics of the data used in these studies, this is due mostly to the intrinsically small angular size of nuclear rings, and the dusty environments in which they exist, both of which make rings non-trivial to observe with most common instruments yielding kinematic data. In this paper we endeavor to  start rectifying this situation by presenting a more homogeneous data set on thirteen nuclear rings.  The nuclear rings studied spectroscopically in the current paper were all part of the sample of \\cite{M08} (hereafter known as M08), who provided photometric insight into the strong star forming nature of nuclear rings. From their \\halpha\\ observational study of 22 nuclear rings, they find that many of the rings exhibit a well-defined age distribution pattern. M08 conclude that this occurrence can be a result of the combination of bar-induced dynamics and gravitational instabilities which are occurring in the proximity of the rings. These results add credence to the idea that the mass inflow driven by the bar along the dust lanes to the ring contact points is a key requirement for the fueling of active (i.e., star forming) nuclear rings (and might also yield a possible signature in the form of non-circular motions). Further evidence for this idea from detailed observations of a handful of galaxies has been presented by, e.g., \\cite{A06, B08}, and \\cite{JK10}.  We now use kinematic information within and just outside the nuclear rings in 13 galaxies to investigate the possible relations between the optical morphology of the rings, the parameters derived from the rotation curve, and their resonant nature. We begin by presenting an overview of the observations and morphological properties of the \\halpha\\ sample in Section~2. We discuss the basic data reduction method, which includes the construction of the velocity field maps, in Section~3. In Section~4 we discuss the kinematic parameters associated with the ring rotation curves, and present results pertaining to non-circular motions in and around the nuclear rings, including velocity residuals. Our detailed analysis of the results is discussed in Section~5, with concluding remarks in Section~6. ",
        "conclusions": "We combined DensePak integral field unit and TAURUS Fabry-Perot observations for a sample of 13 galaxies that contain nuclear rings. Based on this sample, we note the following new findings:  \\begin{itemize} \\item{Nuclear rings are intrinsically circular and unaffected by the local non-axisymmetric environment that can exist from active star formation within the rings; they are influenced by global phenomenon that can affect the potential, such as a bar.} \\item{The compactness, derived from the rotation curves, decreases as the ring width increases.} \\item{The compactness also decreases as the relative ring size increases} \\item{As the relative nuclear ring size increases so does the width of the ring} \\item{Strong bars with a $Q_{b}$ of 0.4 or greater can only be accompanied by small and thin nuclear rings. For smaller torques, the ring size can vary.} \\end{itemize}  We have also reaffirmed several standing findings: \\begin{itemize} \\item{True to resonance theory, the location of the nuclear rings correlates well with the turnover radius.} \\item{The highest torque values occur in those galaxies with the smaller rings relative to the size of their host disk.} \\end{itemize}  In those cases with sufficient radial coverage, we agree that \\begin{itemize} \\item{For NGC~4314 residual velocities sharply decrease and plateau within the nuclear rings.  Ordered velocity patterns increase near the outer edge of the ring (both approaching and receding) close to both sides of the bar minor axis.} \\item{The residual velocities for NGC~5248 near the western side of the bar minor axis reach 40$\\kms$, with a weaker-defined, although evident, transition zone as radii approach the outer ring boundary.} \\item{In the case of NGC~1530 large velocity excesses exist to the west of the nuclear ring, near its exterior edge and then sharply decrease over a small transition zone near the outer boundary of the nuclear ring.} \\end{itemize}"
    },
    "1107/1107.0441_arXiv.txt": {
        "abstract": "{ We present high angular resolution spectral line observations of ground state OH maser emission towards the young high mass proto-stellar object \\object{IRAS 19092+0841}. The two OH main spectral lines, 1665 and 1667 MHz were observed using MERLIN. The two OH lines have been detected close to the central object. The interferometric data have been reduced and a map for both observed frequencies have been produced. The positions and velocities of the OH maser components have been determined. The polarization properties of the OH maser components were determined as well. The main components OH maser components are spread over a region of $\\sim5''$ corresponding to 22400 AU (or $\\sim$ 0.1 pc) at a distance 4.48 kpc. The spread of the relatively small velocity range of the OH masers over a relatively wide distribution and their association of the 44 GHz class I methanol masers indicate that the OH masers emission is probably associated with the external shell of a collapsing core. With the absence of any sign of a disk, outflow or an UCHII region, this would be the first OH masers associated with a star forming region in such an early evolutionary stage. High angular resolution observations of other tracers are needed to reveal the nature of this source. The association between OH masers and the two classes of methanol masers could be a crude evolutionary indicator, such that an association OH masers-class I methanol masers would indicate an earlier evolutionary state than the association OH masers-class II methanol masers. Surveys of class I and class II methanol masers towards OH masers sources are needed to confirm this suggestion. ",
        "introduction": "OH masers observations show that they are associated with different evolutionary stages. It was believed that they are associated with only HII regions. However observations which have been carried out towards of a number of star forming regions show that OH masers are also associated with an earlier stage before the appearance of ionized HII regions (e.g. Cohen et al \\citealp{cohen88}; Braz et al. \\citealp{braz90}). This type of OH maser is associated with an accretion phase, outflow, and circumstellar disks (Brebner et al. \\citealp{brebner87}; Hutawarakorn, \\& Cohen \\citealp{hutawarakorn99}; Hutawarakorn et al \\citealp{hutawarakorn02}; Fuller et al \\citealp{fuller01}; Edris et al. \\citealp{edris05}). An OH maser survey by Edris, Fuller, \\& Cohen (\\citealp{edris07}, hereafter EFC07) detected the OH masers towards 26\\% towards a sample of 217 High Mass Protostellar Objects (HMPOs) candidates. It is of interest to know the distributions of the OH masers and their associations towards these regions in addition to their position in the evolutionary sequence. Maser emission gives a unique opportunity of observing those typically far regions in some details. This paper is the first follow up of EFC07 survey with a high angular resolution observation. One of EFC07 objects (IRAS 20126+4104, hereafter IRAS20126) was studied by Edris et al. (\\citealp{edris05}) at the three masers types, OH, \\htwoo\\, and class II \\methanol\\ masers.  The present object of study, IRAS~19092+0841 (hereafter IRAS19092), is one of the HMPOs sample studied by Molinari et al. (\\citealp{molinari96}; \\citealp{molinari98}; \\citealp {molinari00}; \\citealp{molinari02}). Molinari et al (1996) so-called '\\textit{high} \\& \\textit{Low}' sources are a sub-sample from the Palla et al (1991) sample of 260 IRAS sources (basically candidates for compact molecular cloud association as defined by Richards et al. (1987)). \\textit{High} sources are the one located in the 'higher' part of the colour-colour diagram such that: [25-12]>=0.57 (which in essence is the prescription for the presence of associated UCHII regions according to Wood \\& Churchwell 1989 criteria). On the other hand, '\\textit{Low}' sources are most probably made of 2 distinct evolutionary groups: a very 'early' one (before the creation of a 'real' YSO at the centre and consequently before the appearance of any measurable/detectable UCHII regions; and a more evolved group (older than the '\\textit{High} sources group' containing objects which have already dispersed much of their circumstellar material. And, IRAS19092 belongs to the latter category. % IRAS19092 is at an estimated distance of 4.48 kpc (Molinari et al. \\citealp{molinari96}) and has a luminosity of 9200 \\Lsun\\ (Molinari et al. \\citealp{molinari02}). An ammonia core was found towards IRAS19092 at velocity $\\sim$ 58 km~\\s\\ (Molinari et al. 1996). Although IRAS19092 is a \\textit{Low} type Object, it is one of the few sources associated with different types of maser emission.  Water maser was detected by Brand et al. (1994) and Palla et al. (\\citealp{palla91}). Class II and class I \\methanol\\ masers were detected by Szymczak et al. (\\citealp{szymczak00}), Kurtz et al. (\\citealp{kurtz04}) and recently Fontani et al (\\citealp{fontani10}). OH maser was originally reported by MacLeod et al. (1998). CO and H$_2$ observations failed to detect any outflows from this region (Zhang et al. \\citealp{zhang05}; Varricatt et al. \\citealp{varricatt10}). It is not associated with a close radio continuum emission as well. The closest radio emission detected by Molinari et al. (\\citealp{molinari98}) at 6 cm is offset from the IRAS position by 110 arcsec.  To determine what the OH masers trace and how they are distributed and related to other tracers, IRAS19092 have been observed at high angular resolution using MERLIN. The details of the observations and reduction are given in Sec. 2 and the results presented in Sec. 3. In Sec. 4 we discuss the interpretation while conclusions are drawn in Sec. 5. \\begin{table} \\centering \\begin{tabular}{|c|c|}  \\hline {Observational parameters}  & {OH masers}  \\\\ \\hline Date of observation & 4  \\& 5 April 2003     \\\\ Antenna Used & seven antenna  \\\\ Field centre (2000) & $\\alpha  19^{h} 11^{m} 37.40^{s}$ \\\\ & $\\delta  08^{\\circ} 46' 30.00''$  \\\\ Rest frequency (MHz) & 1665.402 \\\\ & 1667.359  \\\\ No. of frequency channels & 512  \\\\ Total bandwidth (MHz) & 0.5 \\\\ bandpass calibrator & 3C84  \\\\ Polarization angle calibrator & 3C286  \\\\ Phase calibrator & 1919+086 \\\\ \\hline \\end{tabular} \\caption{Observing and calibration parameters for the MERLIN spectral-line observations of IRAS~19092+0841} \\end{table} ",
        "conclusions": "Ground state OH maser emission at 1665 and 1667~MHz was observed towards IRAS19092 at high angular resolution using MERLIN. The OH maser spots are spread over a region of $\\sim5''$ corresponding to 22400 AU (or $\\sim$ 0.1 pc) at a distance 4.48 kpc. We identify one Zeeman pair of OH masers and the splitting of this pair indicates a magnetic field of strength $\\sim4.3$ mG. The velocity range of the OH maser emission suggests that this source is in a collapsing state and the masers arise from the external shell of the collapsing core. With the absence of any sign of disk, outflow or HII region, this would be the first OH masers to be associated with a star forming region in such an early evolutionary stage.  The results of this paper suggest a close association between the OH masers and 44-GHz class I methanol masers in early evolutionary stage. The association between the OH masers and the two classes of the methanol masers could be used as a crude evolutionary indicator with the close association between the OH masers and class I indicates an earlier evolutionary state than its association with class II.   \\vspace{1cm}  ACKNOWLEDGMENTS  MERLIN is a national facility operated by the University of Manchester at Jodrell Bank Observatory on behalf of PPARC."
    },
    "1107/1107.5863_arXiv.txt": {
        "abstract": "We have developed a new method(E-HOLICs) of estimating gravitational shear by adopting an elliptical weight function to measure background galaxy images in our previous paper. Following to the previous paper where isotropic Point Spread Function(PSF) correction is calculated, we consider an anisotropic PSF correction in this paper in order to apply E-HOLICs for real data. A an example, E-HOLICs is applied to Subaru data of massive and compact galaxy clusters A370, and is able to detect double peaks in the central region of the cluster consistent with the analysis of strong lensing. We also study the systematic error in E-HOLICs using STEP II simulation. In particular we consider the dependences of signal to noise ratio \"S/N\" of background galaxies in the shear estimation. Although E-HOLICs does improve systematic error due to the ellipticity dependence as shown in part 1, a systematic error due to the S/N dependence remains, namely E-HOLICs underestimates shear when background galaxies with low S/N objects are used. We discuss a possible improvement of the S/N dependence. ",
        "introduction": "Weak gravitational lensing has been widely recognized as a unique and very powerful method to study not only mass distribution of the universe but also the cosmological parameters (see for example, Mellier 1999, van Waerbeke \\& Mellier 2003, Schneider 2006, Munshi et al. 2008). In particular, cosmic shear has attracted much attention recently because of its potential to determine the so-called cosmic equation of state, namely the relation between the effective energy density and pressure of the dark energy. There are some detections of cosmic shear(Bacon et al 2000; Kaiser et al 2000; Maoli et al 2001; van Waerbeke et al 2001: Refregier ey al 2002; Bacon et al 2003; hamana et al 2003, Casertano et al 2003; van Waerbeke et al 2005; Massey et al 2005; Hoekstra et al 2006). However the lensing signal of cosmic shear is very small and thus highly accurate shear measurement is required. The most popular method of shear estimation is given by Kaiser et al. 1995(called the KSB method: see also Luppino \\& Kaiser 1997; Hoekstra et al. 1998; Viola et al. 2011) where change in the moments of galactic light distribution by lensing is extracted from the measurement. Other methods of shear estimation have been also developed (Bernstein \\& Jarvis 2002; Refregier 2003; Kuijken et al. 2006; Miller et al. 2007; Kitching et al. 2008; Melchior 2011). In relation with our E-HOLICs method, we mention the work by Bernstein \\& Jarvis 2002 and Melchior 2011 who also introduced an elliptical weight for the shape measurement. All of them are very challenging and successful for cluster lensing. Although some of them have already been used for cosmic shear, it is argued that none of them have achieved the required level of accuracy to measure the cosmic equation of state in a few percent level.  On the other hand there are planned cosmic shear observations in very near future for the purpose of measuring the cosmic equation of state. In this situation, it is urgent to develop sufficiently accurate method of shear estimation. In our previous paper(Okura \\& Futamase 2011, we call part 1) we have developed a new method of the evaluating shear which is based on KSB method by adopting an elliptical weight function to measure the shape of background galaxies. Our method is a natural development of our previous studies of weak lensing analysis which uses higher order multiple moments of shape of background galaxies. We called the method as HOLICs(Higher Order Lensing Image Characteristics) which are quantitates with a definite spin properties made of higher-order multipole moments(Okura et al. 2007). We have shown that Oct-HOLICs is an unbiased measure of Flexion and is successfully applied to a galaxy cluster Abell 1689 to reveal substructure in the central part of the cluster(Okura et al. 2008) and spin-2 HOLICs which is spin-2 combination of higher order moments increases the lensing information of image and is a useful method to apply cosmic shear measurement(Okura \\& Futamase 2009). We called our method as generalized as elliptical weighted HOLICs(E-HOLICs) method because of the elliptical weight. We have showed that HOLICs method with an elliptical weight is able to measures the lensing distortion more accurately by weighting highly to brighter region of image than in the standard KSB method, and thus it can reduce effects of systematic error and random noise more effectively than KSB method. We have also calculated isotropic PSF correction in part 1 and tested its accuracy using STEP 2 simulation data. Following to part 1, we present a method of correcting PSF containing anisotropic part in this paper. This is necessary to apply the method to real data. In the case of isotropic PSF correction, it was shown that the corrected shear can be obtained by solving a polynomial equation.  However, since anisotropic part has an arbitrary direction, it is natural to expect that the shear after PSF correction is obtained by solving coupled two polynomial equations. We succeed in reducing these coupled equations into tractable forms by dividing anisotropic part of PSF into parallel part and orthogonal part against the direction of shear(see section 2 in detail). We test E-HOLICs method with STEP2 simulation data. In part 1, we can only test the data with isotropic PSF, but in this paper we test all PSF set and obtain more detailed results. After establishing PSF correction, we are able to apply our method to real data of galaxy cluster A370.  The organization of the paper is as follows. First we give a brief review of weak lensing and E-HOLICs method in section 2. Then in section 3, we describe the correction of PSF with anisotropic part in E-HOLICs method by dividing anisotropic part of PSF into anisotropic part into parallel part and orthogonal part. In section 4 we tested E-HOLICs method by using STEP2 simulation data, in particular we investigate the dependences of S/N and size of background galaxies in the shear estimation. Then we apply E-HOLICs method to real data of galaxy cluster A 370. It will be shown that E-HOLICs is able to detect two peaks in the central part of the cluster consistent with strong lensing analysis. Finally we summarize our results and give some discussion in section 6. ",
        "conclusions": "It is widely recognized now that there are many sources of systematic errors in the evaluation of weak lensing shear and their improvement is urgently required for the precise measurement of cosmic shear. We have developed a new method of shear estimation, based on the KSB method and called \"E-HOLICs method\", by introducing the elliptical weight function to define multipole moments of galaxy light distribution in  previous paper(part 1). The use of elliptical weight function is expected to avoid the systematic error coming from an expansion of the weight function, which was usually done in some of the previous approaches to shape measurement, including the KSB method. Bernstein \\& Jarvis 2002 and Melchior 2011 are also developing schemes which use an elliptical weight for the shape measurement.   Following part 1 we developed the correction scheme of PSF with anisotropy in E-HOLICs method. Generally, the direction of the elliptical PSF is not same with the direction of distortion, and thus we must treat these two direction in PSF correction, By dividing anisotropic part of PSF into parallel and orthogonal part against the direction of the distortion, we succeeded to reduce PSF correction into two steps of simple one dimensional form. Then we test E-HOLICs using STEP 2 data simulation. In particular we tested the dependence on the S/N and size of background galaxies. We have found that E-HOLICs gives underestimation for low S/N sources. Although the precision is of the same level of other methods, this causes a serious problem in the accurate shear measurement. We have found that at least a part of reason for the underestimation is the errors in the centroid. Although the error in the position of centroid distributes randomly, the area in the image where the measured ellipticity from the centroid within it is smaller than the correct value is larger than the other area in the image. Thus the random error in the centroid has a tendency to cause the underestimation of the measured shear.  We are investigating this effect and the result will be shown in the forthcoming publication.  In this paper we have ignored higher-order moments in PSF which might be important in some cases. We aim to study the higher-order PSF correction also in the future publication.  As we pointed out, E-HOLICs method requires longer time for measuring moments if we use many stars for PSF correction than other methods such as KSB method. This might cause a practical difficulty in any wide field surveys. We hope that there will be a way to avoid this problem by using powerful computer specially programed for this purpose.    As an application of E-HOLICs methods, we analyzed a massive and compact galaxy cluster Abell 370 using by Subaru/S-Cam data. We found that E-HOLICs can detect two mass peaks which is consistent with the strong lensing analysis. The KSB method cannot detect two peaks possibly because available number of background objects is reduced in compared with E-HOLICs.  Although the E-HOLICs method developed in this paper has a potential to accurately measure the shear, it has still some shortcoming as described above and definitely more studies are necessary to apply E-HOLICs to the planned large-scale cosmic shear observations.  We thank K. Umetsu for a useful discussion and for providing his scripts, N. Okabe for useful discussion. This work is supported in part by the GCOE Program \"Weaving Science Web beyond Particle-Matter Hierarchy\" at Tohoku University and by a Grant-in-Aid for Scientific Research from JSPS(nos. 18072001, 20540245 for TF), as well as by Core-to-Core Program \"International Research Network for Dark Energy.\"  \\appendix"
    },
    "1107/1107.2391_arXiv.txt": {
        "abstract": "{We review the spectral properties of the black hole candidate Cygnus~X-1.  Specifically, we discuss two recent sets of multi-satellite observations.  One comprises a 0.5--500\\,keV spectrum, obtained with \\emph{every} flying X-ray satellite at that time, that is among the hardest \\cyg\\ spectra observed to date.  The second set is comprised of 0.5--40\\,keV \\chandra-\\hetg\\ plus \\rxte-\\pca\\ spectra from a radio-quiet, spectrally soft state.  We first discuss the ``messy astrophysics'' often neglected in the study of \\cyg, i.e., ionized absorption from the wind of the secondary and the foreground dust scattering halo. We then discuss components common to both state extremes: a low temperature accretion disk, and a relativistically broadened Fe line and reflection.  Hard state spectral models indicate that the disk inner edge does \\emph{not} extend beyond $\\gtrsim 40\\,GM/c^2$, and may even approach as close as $\\approx 6\\,GM/c^2$. The soft state exhibits a much more prominent disk component; however, its very low normalization plausibly indicates a spinning black hole in the \\cyg\\ system. ",
        "introduction": "After the initial discovery \\citep{bowyer:65a} of the black hole candidate \\cyg, \\citet{tananbaum:72a} noted that, at least in the 2--10\\,keV band, it exhibited two distinct spectral states: a bright, spectrally soft state, and a fainter, spectrally hard state. Subsequent broad-band studies further elucidated this dichotomy.  The soft state spectrum peaks between 1--2\\,keV, has a weak hard tail with photon spectral index $\\Gamma > 2.1$, and is radio quiet.  The hard state spectrum is well-described by an exponentially cutoff (folding energy 125--250\\,keV) broken powerlaw with 2-10\\,keV photon index $\\Gamma < 2.1$ and a harder photon index at $\\gtrsim 10$\\,keV, weak soft excess, and it is radio loud.  \\cyg\\ spends the majority of its time in its hard state.  A range of such spectra (and further references) can be found in the work of \\citet{wilms:06a}, which describes a multi-year radio/\\rxte\\ monitoring campaign of \\cyg.  \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{nowak_2011_01_fig01.ps}} \\caption{ \\footnotesize \\cyg\\ in its two most extreme spectral states: the hard state dominated by a cutoff powerlaw spectrum extending to $> 100$\\,keV, and the soft state dominated by a thermal spectrum peaking between 1--2\\,keV.  The above hard state is from a multi-wavelength campaign wherein \\cyg\\ was observed by \\emph{every} flying X-ray satellite.  Here we display \\chandra-\\hetg\\ spectra (histogram), \\suzaku-\\xis\\ (solid diamonds) and -\\gso\\ (hollow triangles) spectra, \\rxte-\\pca\\ (circles) and -\\hexte\\ (hollow diamonds) spectra, and \\integral-\\spi\\ (solid triangles) spectra (see \\protect{\\citealt{nowak:11a}}).  The soft state spectra are from a simultaneous \\chandra-\\hetg\\ and \\rxte\\ campaign conducted in January 2011. The above spectra are presented without reference to any underlying spectral model, and are unfolded using only the spectral response of the detectors (see \\protect{\\citealt{nowak:05a}}).} \\label{fig:states} \\end{figure}  Two extreme examples of the states are presented in Fig.~\\ref{fig:states}.  This figure shows April 2008 observations of a spectrally hard state, covering the bandpass of 0.5--500\\,keV.  It is among the hardest spectra seen from \\cyg; a description of these spectra can be found in \\citet{nowak:11a}.  Also shown are spectra from simultaneous \\chandra-\\hetg/\\rxte-\\pca\\ observations, covering the 0.5--40\\,keV range of a spectrally soft state.  In terms of \\emph{observed} flux, the 0.1--50\\,keV flux of the soft state is $3.9 \\times 10^{-8}\\,{\\rm erg\\,cm^{-2}\\,s^{-1}}$, while the \\emph{observed} 0.5--300\\,keV flux of the hard state is $4.9 \\times 10^{-8}\\,{\\rm erg\\,cm^{-2}\\,s^{-1}}$.  However, making plausible corrections for absorption and extrapolating the spectra to determine bolometric luminosities, this is at 2.2\\%\\,${\\rm L}_{\\rm Edd}$ and this hard state is at 1.6\\%\\,${\\rm L}_{\\rm Edd}$.  (The Eddington luminosity, ${\\rm L_{\\rm Edd}}$, used here assumes a distance of 1.86\\,kpc and a mass of 15\\,${\\rm M}_\\odot$; \\citealt{reid:11a,orosz:11a}.)  That is, whereas we directly observe most of the flux in the hard state, nearly 2/3 of the soft state flux is \\emph{unobserved} due to both absorption and bandpass limitations.   \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true,bb=30 340 550 770]{nowak_2011_01_fig02.ps}} \\caption{ \\footnotesize A typical hardness-intensity diagram for a black hole transient outburst (based upon an outburst of GX~339$-$4; \\protect{\\citealt{belloni:05a,homan:05a}}).  Transients begin faint, hard, and radio loud; they evolve to brighter states while remaining hard; they soften and become radio quiet (often preceded by a radio ejection event); they then fade, harden, and become radio loud once more.  \\cyg, which is persistently emitting in the X-ray band, occupies only a small portion of this diagram at fractional Eddington luminosities of $\\approx 2\\%$, near the so-called radio loud/radio quiet transition ``jet line''.} \\label{fig:q} \\end{figure}  As has been noted by previous researchers, the range of bolometric luminosities traversed by \\cyg\\ spans only a factor of $\\approx 3$--4 \\citep[][and references therein]{wilms:06a}.  This is somewhat narrow compared to most black hole transients.  Furthermore, \\cyg\\ also traverses a narrower range of colors than many black holes, never exhibiting a purely disk-dominated spectrum without a hard tail \\citep[e.g., like the simple disk-dominated spectrum of 4U~1957+11;][]{nowak:08a}.  Black hole transients often follow color-intensity diagrams as shown in Fig.~\\ref{fig:q}, spanning $\\gtrsim 3$ orders of magnitude in luminosity, and wider extremes of color variations.  \\cyg, however, exists on an especially interesting portion of this `q-diagram' --- moving between the radio quiet/soft $\\Leftrightarrow$ radio loud/hard-intermediate state transition near $\\approx 2$\\%\\,${\\rm L}_{\\rm Edd}$.  \\begin{figure}[] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{nowak_2011_01_fig03.ps}} \\caption{ \\footnotesize Simultaneous \\suzaku-\\xis\\ (diamonds) and \\chandra-\\hetg\\ (histogram) spectra fit with a simple phenomenological model consisting of a disk, powerlaw, a relativistically broadened line, and a narrow line, modified by ionized absorption (see \\protect{\\citealt{hanke:08a}} and \\protect{\\citealt{nowak:11a}}). Furthermore, the \\hetg\\ spectra are modified by dust scattering.  One set of models, and all of the residuals, are shown with both the ionized absorption and dust scattering absent.}\\label{fig:messy} \\end{figure}  The hypothesized emission components in the \\cyg\\ system include an accretion disk, a Comptonizing corona (with either a thermal, or hybrid thermal/non-thermal electron distribution; \\citealt{coppi:99a}) that upscatters photons from the disk, a radio emitting jet that might also contribute to the X-ray \\citep{markoff:05a}, and relativistically smeared reflection/fluorescence from the disk.  The important questions are, what are the relative contributions of these components, and how do they and the system geometry change between the two state extremes?  \\begin{figure}[ht] \\resizebox{0.86\\hsize}{!}{\\includegraphics[clip=true]{nowak_2011_01_fig04a.ps}} \\resizebox{0.86\\hsize}{!}{\\includegraphics[clip=true]{nowak_2011_01_fig04b.ps}} \\resizebox{0.86\\hsize}{!}{\\includegraphics[clip=true]{nowak_2011_01_fig04c.ps}} \\caption{\\footnotesize \\cyg\\ hard state \\suzaku/\\rxte\\ spectra, fit with different models. (Model components shown individually: seed photons, Compton spectrum, synchrotron, SSC, disk, reflection, broad and narrow lines.)  Top: thermal Comptonization with a high seed photon temperature.  Middle: non-thermal Comptonization with a low seed photon temperature.  Bottom: jet model dominated by synchrotron and SSC emission.  All models also have: a low temperature disk, relativistically broadened reflection and Fe fluorescence line, and a narrow Fe line.  Disk and broad Fe line parameters are consistent with the disk extending very close to the innermost stable circular orbit \\protect{\\citep{nowak:11a}}.} \\label{fig:hard} \\end{figure} ",
        "conclusions": ""
    },
    "1107/1107.5264_arXiv.txt": {
        "abstract": "We propose, in a heuristic way, a relativistic modified gravity model as an alternative to particle dark matter at galactic scales. The model is based on a postulated preferred time foliation described by a dynamical scalar field called the ``Khronon''. In coordinates adapted to the foliation it appears as a modification of general relativity violating local Lorentz invariance in a regime of weak gravitational fields. The model is a particular case of non-canonical Einstein-{\\ae}ther theory, but in which the {\\ae}ther vector field is hypersurface orthogonal. We show that this model recovers the phenomenology of the modified Newtonian dynamics (MOND) in the non-relativistic limit, and predicts the same gravitational lensing as general relativity but with a modified Poisson-type potential. ",
        "introduction": "\\label{secI}  Modifying gravity in the regime of weak gravitational fields or accelerations, is the leitmotiv of MOND --- the Modified Newtonian Dynamics \\cite{Milg1,Milg2,Milg3}. The MOND paradigm is very successful at solving the problem of the flat rotation curves of galaxies without the need of dark matter (see \\cite{SandMcG02} for a review), and at explaining the correlations between the distribution and dynamics of ordinary matter versus dark matter \\cite{TF,McG11}. However, it is based on a non-relativistic formula and needs to be generalized in order to be applied in cosmology and to the gravitational lensing observed at galaxy cluster scale \\cite{GerbalDurret,Clowe,ASZF07}, which are part of the great successes of the standard cosmological paradigm $\\Lambda$-CDM based on particle cold dark matter \\cite{BHS05} and the cosmological constant $\\Lambda$.  We adopt for MOND the modification of the non-relativistic Poisson equation~\\cite{BekM84}, \\begin{equation}\\label{MOND} \\vecnab \\cdot \\left[ \\mu\\left(\\frac{g}{a_0}\\right) \\vecnab \\phi \\right] = 4\\pi G \\rho\\;, \\end{equation} where $\\rho$ is the density of ordinary (essentially baryonic) matter, $\\phi$ is the gravitational potential, $\\bm{g}=-\\vecnab \\phi$ is the gravitational field and $g = \\vert\\bm{g}\\vert$ its ordinary Euclidean norm. The MOND function $\\mu$ depends on the single argument $g/a_0$, where $a_0=1.2\\times 10^{-10}\\,\\mathrm{m}/\\mathrm{s}^2$ is a constant acceleration scale. In the limit of weak gravitational fields \\textit{i.e.} $g\\ll a_0$, the MOND function is linear, $\\mu(g/a_0)=g/a_0$, while in the strong field regime $g\\gg a_0$ (though non relativistic), $\\mu$ tends to one so that \\eqref{MOND} reduces to the usual Poisson equation.  Several relativistic extensions of MOND have been proposed (see \\cite{BGef07,CFPS11} for reviews). The Tensor-Vector-Scalar (TeVeS) theory of Bekenstein and Sanders \\cite{Sand97,Bek04,Sand05} extends general relativity with a time-like vector field and one scalar field. Einstein-{\\ae}ther theories \\cite{JM00,JM04} involve a unit time-like vector field which is non-minimally coupled to the metric, and were originally used to describe the phenomenology of local Lorentz invariance violation. Those theories, when modified to involve non-canonical kinetic terms, have been shown to provide interesting examples of relativistic MOND theories \\cite{ZFS07,HZL08}. The vector field postulated in TeVeS is in fact analogous to the vector field in Einstein-{\\ae}ther theories \\cite{ZFS06,CWW08}. The cosmology in TeVeS and Einstein-{\\ae}ther theories has been extensively studied \\cite{SMFB06,LMB07,Sk08,ZZ10}. The parametrized post-Newtonian parameters, included the preferred frame parameters crucial in the presence of a vector field, have been worked out for both the canonical Einstein-{\\ae}ther theory \\cite{FJ06} and TeVeS \\cite{Sagi09}. More recent proposals for relativistic MOND include a bimetric theory of gravity \\cite{Milgbim1,Milgbim2}, a refinement of TeVeS using a Galileon ``$k$-mouflage'' to prevent deviations from general relativity at small distances \\cite{BDEF11}, and a non-local extension of general relativity \\cite{DEFW11}. One can also invoke some new exotic properties of the dark matter rather than modifying gravity \\cite{BL08,BL09}.  On the other hand, much interest focused recently on the possible violation of the local Lorentz invariance in a completion of general relativity at high energy. This Ho\\v{r}ava-Lifshitz approach \\cite{Horava09} is motivated by the construction of a power-counting renormalizable quantum theory of gravity. The Lorentz invariance violation results from a preferred time foliation which is described in a $4$-dimensional covariant formulation by a dynamical scalar field sometimes called the ``Khronon'' \\cite{BPS11}. A consistent extension of the (non-projectable version of the) Ho\\v{r}ava gravity has been proposed \\cite{BPS10,BPS11} to avoid some problems regarding the stability of the theory. In this extension the added terms crucially involve the acceleration of the congruence normal to the preferred family of hypersurfaces. It has been shown that the low energy limit of the extended Ho\\v{r}ava gravity is equivalent to an Einstein-{\\ae}ther theory in which the vector field is hypersurface-orthogonal \\cite{Jacobson10,BPS11}.  In the present paper, we revisit some works on Einstein-{\\ae}ther theories with non-canonical kinetic terms (notably \\cite{ZFS07,HZL08}), but in the framework of a preferred time foliation, inspired by these recent approaches to quantum gravity. We propose a relativistic MOND theory which modifies general relativity by terms depending on the acceleration in a way somewhat similar to the extended Ho\\v{r}ava gravity \\cite{Horava09,BPS10,BPS11}. However, our motivation is quite different from that in \\cite{Horava09,BPS10,BPS11}, as it is purely classical, and we are not concerned with the problem of quantizing gravity. In particular, the violation of Lorentz invariance will be effective at low energy, for very weak gravitational fields. Our aim is to provide a phenomenologically viable alternative to the particle cold dark matter (at least at galactic scales).  The modification we propose consists of introducing in the Lagrangian of general relativity a single free function of the norm $a$ of the acceleration of the congruence normal to the preferred hypersurfaces. (Unlike in \\cite{BPS10,BPS11}, we are not considering a finite number of terms including this acceleration, each one with a coupling constant scaling with the Planck mass, but a fully non-linear function.) The theory will be a particular case of non-canonical Einstein-{\\ae}ther theory but in which the {\\ae}ther field is hypersurface orthogonal. We shall present the theory in both a 3+1 decomposition and $4$-dimensional covariant form, and shall explicitly prove the equivalence of the two formulations. In agreement with general studies on non-canonical Einstein-{\\ae}ther theories \\cite{ZFS07,HZL08}, we will show that it is possible, by specifying the asymptotic behaviour of the function of the acceleration when $a\\ll a_0$, to recover the Bekenstein-Milgrom equation \\eqref{MOND} in the non-relativistic limit. The theory will also be viable regarding the gravitational lensing of photons, which will be the same as in general relativity but with a Newtonian potential obeying the modified Poisson equation.\\footnote{While this work was in preparation, a paper based on similar ideas appeared on the archives \\cite{Sand11}. However, the field equations in that paper are in disagreement with ours.}  The plan of this paper is as follows. In Section \\ref{secII} we present the $4$-dimensional version of the theory, introducing the dynamical Khronon field defining the space-time foliation, and investigating the resulting field equations. In Section \\ref{secIII} we switch to an alternative 3+1 point of view, for which the time coordinate coincides with the Khronon field, and where the new terms associated with the acceleration in the action may be interpreted as Lorentz invariance breaking terms due to a preferred-frame effect. In Section \\ref{secIV}, we investigate the non-relativistic limit and the conditions under which one can retrieve the non-relativistic MOND equation \\eqref{MOND}. Section \\ref{secV} contains some concluding remarks. Finally, Appendix \\ref{appA} is devoted to the proof of equivalence between the $4d$ covariant formulation and the 3+1 formalism, while Appendix \\ref{appB} presents the canonical decomposition of the energy-momentum tensor associated with the Khronon field when interpreted as a matter field. ",
        "conclusions": "\\label{secV}  We have proposed a relativistic modified gravity model, based on a preferred 3+1 space-time foliation, reproducing the phenomenology of MOND \\cite{SandMcG02} in the weak-field limit. The modification with respect to general relativity consists in adding to its ordinary Lagrangian a function of the norm of the acceleration of the congruence associated with this foliation. We investigated two different points of view on this theory: In the first, we introduced a scalar field called the Khronon defining the foliation by its constant-value hypersurfaces, keeping a full $4$-dimensional covariant formalism, while in the second, we wrote the theory in a 3+1 fashion, in a frame where the time coordinate coincides with the Khronon field and where the Lagrangian is no longer manifestly Lorentz invariant. The spirit of our approach is similar as that of recent attempts at building a consistent quantum theory of gravity \\cite{Horava09,BPS10,BPS11}, but here the Lorentz invariance violation occurs at low energy, and our motivation is purely classical. We showed the equivalence between the $3d$ and $4d$ formulations of the model, and gave the requirements on the function initially introduced in the Lagrangian to recover MOND in the non-relativistic approximation.  We leave to future work some important questions. First, it is easy to see that in an homogeneous and isotropic background, our model simply reduces to the addition of a cosmological constant. However the cosmological implications of this model in a perturbative regime have not been studied yet. We have in mind the viability of the model when faced to cosmological observations, notably the anisotropies of the cosmic microwave background and the structure formation (see \\cite{SMFB06,LMB07,Sk08,ZZ10} for general investigations of the cosmology in TeVeS and Einstein-{\\ae}ther theories).  A second question is that of the viability regarding Solar system tests and binary pulsars data. Although it should recover general relativity plus a cosmological constant in the strong field regime appropriate to these systems, see \\eqref{lambdaGR},\\footnote{For instance, the MOND transition radius for the Sun is $\\sqrt{G M_\\odot/a_{0}} \\simeq 7100 \\text{AU}$.} our model includes a preferred frame effect which we would like to quantify precisely, using various hypothesis for the motion of the Solar system with respect to the preferred frame. (See \\cite{FJ06,Sagi09} for the computation of the preferred frame post-Newtonian parameters in TeVeS and canonical Einstein-{\\ae}ther.) More importantly, the preferred frame effect should also be computed in the MOND regime, using various forms for the MOND function, to see how it affects the usual MOND fit of the flat rotation curves of galaxies; see \\eqref{aNmod} and the discussion right after.  Another question is related to the dynamics of the Khronon field itself [obeying the equation \\eqref{Teq}]. Indeed in this model the foliation is dynamical, and the evolution of the Khronon field should be compatible with the preservation of its smoothness and regularity properties allowing a space-time foliation to be built on it."
    },
    "1107/1107.2819_arXiv.txt": {
        "abstract": "We review recent progress in numerical relativity simulations of black-hole (BH) spacetimes. Following a brief summary of the methods employed in the modeling, we summarize the key results in three major areas of BH physics: (i) BHs as sources of gravitational waves (GWs), (ii) astrophysical systems involving BHs, and (iii) BHs in high-energy physics. We conclude with a list of the most urgent tasks for numerical relativity in these three areas. ",
        "introduction": "\\label{sec:intro} For almost a century now after its discovery, the Schwarzschild solution describing a static, spherically symmetric vacuum spacetime in Einstein's theory of general relativity (GR) has been a valuable tool for a wealth of theoretical and experimental studies, including in particular weak-field tests of the theory.  For more than half a century, the validity of the solution beyond the Schwarzschild radius $r=2GM/c^2$ and the consequent existence of BHs was thought to be a mathematical artifact of the field equations. This picture has changed dramatically in more recent decades. Crucial events in this sense were the mathematical discovery of the Kerr solution describing rotating BHs \\cite{Kerr1963}, and the astrophysical X-ray and radio observations in the 1960s and 1970s. Cosmic X-ray sources, first detected in 1963 \\cite{Giacconi1963}, were soon associated with compact stellar objects and, in particular, with mass exchange from an ordinary star onto a compact object in binary systems \\cite{Shklovsky1967}. This interpretation received further support when the X-ray source Cyg X-1 was identified with the supergiant BOIb star HD~226868 \\cite{Bolton1972, Webster1972}, which established Cyg X-1 as the first observed X-ray binary.  Over the next few years, the explanation of these systems in terms of accretion onto compact stellar objects in binary systems became widely accepted (see \\cite{VanParadijs1998} for a review).  Measurements of the binary period and radial velocity provide lower limits on the mass of the compact companion. The values obtained for some X-ray binaries comfortably exceed the upper mass limit of about $3~M_{\\odot}$ for neutron stars. Recent investigations of Cyg X-1, for example, predict $8.7\\pm 0.8~M_{\\odot}$ for the compact member \\cite{Shaposhnikov2007}.  Also starting in the early 1960s, optical and radio observations identified ``star like'' objects with surprisingly large redshifts \\cite{Matthews1963, Schmidt1963}.  If interpreted as cosmological in origin, the redshifts implied enormous distances and, accordingly, luminosities.  Over the next decade the cosmological origin of these ``quasi-stellar'' sources (or quasars) became clear, and accretion onto supermassive BHs (SMBHs) \\cite{Salpeter1964, Zeldovich1964} was accepted as the most plausible explanation of their energetics \\cite{Rees1978,Melia:2007vt}. By now, observations of stellar dynamics near the centers of galaxies and iron K$\\alpha$ emission line profiles provide strong evidence that many galaxies harbor BHs at their centers \\cite{Kormendy1995, Richstone1998}. BH masses are generally closely correlated to the bulge velocity dispersion and luminosity \\cite{Ferrarese2000, Gebhardt2000}.  For the Milky Way, astrometric and radial velocity observations of short-period stars orbiting the galactic center \\cite{Schoedel2002, Ghez2008} predict an unseen gravitational mass of $4.1\\pm 0.6 \\times 10^6~M_{\\odot}$.  Electromagnetic observations provide strong but {\\em indirect} evidence for the existence of astrophysical BHs. A more direct way of probing the BH nature of astrophysical compact objects is offered by the newly emerging field of GW astronomy. First-generation ground-based laser-interferometric detectors (such as LIGO, VIRGO and GEO600) have reached design sensitivity, and advanced versions of these detectors will be operational within a few years \\cite{LIGOweb, VIRGOweb, GEO600web}.  One of the most prominent sources of detectable GWs for these detectors is the inspiral and coalescence of stellar mass and, possibly, intermediate-mass BH (IMBH) binaries. At lower frequencies, space-based laser interferometric observations \\cite{ESALISAweb} have the potential to complement ground-based efforts with high signal-to-noise ratio observations of massive BH binaries. Space-based detectors are guaranteed to observe galactic binaries containing white dwarfs and/or neutron stars, and they may also detect the inspiral of stellar compact objects into massive BHs. Due to the weak interaction of GWs with any type of matter, including the detectors, digging physical signals from the noisy data stream represents a daunting task and heavily relies on so-called {\\em matched filtering} techniques \\cite{Finn1992}.  Matched filtering is a common choice to search for signals of known form in noisy data, and works by cross-correlating the actual ``signal'' (i.e., the detector's output) with a set of theoretical templates.  Accurate modeling of the GW sources thus plays a vital role in maximizing the scientific output from these experimental efforts.  Event rate estimates for IMBHs and stellar-mass BH binaries are very uncertain, but advanced Earth-based detectors are expected to observe several events per year (see \\citep{Abadie2010} for a review). Belczynski {\\it et al.} noted that recent metallicity measurements require revisions in the population synthesis models used for event rate estimates, suggesting that stellar-mass BH-BH binaries may be the first GW sources to be detected \\cite{Belczynski2010}.  The observations of two possible X-ray binary precursors of BH-BH binaries provide a much needed observational constraint on compact binaries containing at least one BH, suggesting again that BH-BH binary mergers may be more common than anticipated \\citep{Bulik2011}. Population synthesis models have large uncertainties, related e.g. to the poorly known common envelope phase, and some models would necessarily be ruled out if Advanced LIGO does {\\it not} detect BH binaries. In this sense, GW detectors are {\\em guaranteed} to put constraints on compact binary formation in the near future. Schutz \\cite{Schutz2011} also pointed out that coherent data analysis using three or more detectors may further increase event rates by factors of a few, making prospects of doing astrophysics with Earth-based detectors even brighter. Estimates for event rates of SMBH binaries detectable by space-based interferometers range from $0.1$ to $1000$s per year, with a most likely estimate of $\\sim 20-30$ \\cite{Volonteri2003,Begelman:2006db,Sesana2007,Arun2008a}.  Over the past decade, BH modeling has also attracted a great deal of interest in the context of high-energy physics.  According to the gauge/gravity duality, gravitational physics in anti-de Sitter (AdS) backgrounds describes field theories on the boundary of that spacetime \\cite{Maldacena1997,Witten1998}. According to this duality, BHs in AdS are dual to an equilibrium field theory, and therefore interacting BHs or BHs off-equilibrium may represent interesting physics from the dual point of view. In fact, there is some evidence that some of the physics behind quark-gluon plasma formation in heavy-ion collisions is captured by this duality, which on the gravitational side corresponds to shock wave collisions and BH formation in AdS \\cite{Mateos:2007ay}.  Understanding BH spacetimes is also of interest for high-energy physics in connection with extra-dimensional scenarios. Solutions to the hierarchy problem in physics have been developed where all standard-model interactions are confined to a 3+1 dimensional brane embedded in a spacetime with large extra dimensions or an extra dimension with a warp factor \\cite{Arkani-Hamed1998, Randall1999}. In these models gravity propagates in the entire higher-dimensional spacetime (the ``bulk''), and the fundamental Planck scale could be as low as $1~{\\rm TeV}$. These energy scales are accessible to experiments such as the Large Hadron Collider (LHC).  If the fundamental Planck scale is so low, Thorne's hoop conjecture suggests the exciting possibility of BH production in high-energy parton-parton collisions \\cite{Giddings2002,Dimopoulos2001}, i.e.~a direct observable signature of the existence of extra-dimensions.  Accurate modeling of energy and momentum losses in the form of GWs during the collisions and determination of the BH formation cross section will be crucial for the analysis of data from such experiments (see e.g.~\\cite{Frost2009}), and they should be amenable to classical calculations in $D$-dimensional GR \\cite{Banks1999}.  Finally, the understanding of BH dynamics under extreme conditions is fascinating from a mathematical-physics point of view.  Questions such as BH stability, cosmic censorship, etcetera can only be addressed by solving the full nonlinear set of the field equations.  For many of these scenarios, it is convenient to divide the coalescence of a BH binary in the framework of GR into three stages which mirror the different tools used for their modeling: (i) The inspiral phase, where the interaction between the individual holes is still sufficiently weak to justify the use of post-Newtonian (PN) techniques \\cite{Blanchet2006}; (ii) The final orbits, plunge and merger, which are governed by the strong-field regime of Einstein's equations and can be described only by fully numerical simulations; (iii) The ringdown of the remnant BH, which is amenable to a perturbative treatment \\cite{Teukolsky1973, Leaver1986, Berti2009}.  In this article we will focus on the second, strong-field stage, on the numerical tools dedicated to its study, and results obtained following the 2005 breakthroughs \\cite{Pretorius2005a,Campanelli2006,Baker2006}. As we shall see, however, a comprehensive understanding of BH dynamics requires a close interplay of numerical and analytical studies, and we will discuss in various cases the interface between numerical relativity and approximation techniques.  Following a brief summary of the framework employed in numerical relativity (Sec.~\\ref{sec:numerics}), in Secs.~\\ref{sec:GWs}-\\ref{sec:highenergy} we will present the main results obtained from numerical studies of BHs in the context of GW detection, astrophysics and high-energy physics, respectively. We conclude in Sec.~\\ref{sec:conclusions} with a discussion of future applications of numerical relativity.  At the end of each section we will provide references to the literature for further reading.  \\vspace{3pt} \\noindent {\\bf \\em Notation:} We shall be using geometrical units, such that the gravitational constant $G=1$ and the speed of light $c=1$. We denote spacetime indices $0\\ldots D-1$ by Greek letters, and spatial indices $1\\ldots D-1$ by Latin letters. ",
        "conclusions": "\\label{sec:conclusions}  We conclude this review with a brief summary of the most urgent problems to be tackled by numerical relativity in the three main areas in the near future.  \\vspace{3pt} \\noindent {\\bf \\em Gravitational wave physics:} In order to meet tight accuracy requirements in GW template generation, especially for parameter estimation, we either need longer numerical waveforms or a denser spacing of templates in the parameter space. Both approaches will require computational resources that grow non-linearly as a function of the accuracy, either because of the slow nature of the inspiral at larger BH separations (see Sec.~2 in \\cite{MacDonald2011} for a quantitative discussion) or because of the dimensionality of the parameter space.  A second major task is the extension of existing template models to generic binaries with precessing spins and/or smaller mass ratios, bridging the gap to perturbative modeling of extreme mass ratio binaries.  \\vspace{3pt} \\noindent {\\bf \\em Astrophysics:} While existing formulae for the kick and final spin resulting from the coalescence of BHs have been helpful for astrophysical studies, the calibration of generic models as proposed by Boyle {\\em et al.} \\cite{Boyle2007a,Boyle2007b} requires a more comprehensive exploration of the parameter space. The relatively young field of numerical relativity simulations of BHs surrounded by accretion disks will likely improve our insight into expected optical counterparts to BH binary mergers, and thus provide a vital tool for the area of {\\em multi-messanger astrophysics}.   \\vspace{3pt} \\noindent {\\bf \\em High-energy physics:} Arguably the most urgent task is the determination of the scattering cross-section and the loss of energy and momentum in GWs in trans-Planckian scattering of BHs in {\\em arbitrary} dimensions. In view of the stability problems encountered in present studies, it also appears desirable to obtain a better understanding of the well-posedness of the formulations currently employed for such studies. Second, the modeling of BHs in generic spacetimes such as de Sitter and anti-de Sitter is still in its infancy, even if preliminary studies of BHs in non-asymptotically flat spacetimes are encouraging. A better understanding of boundary issues, well-posedness of the formulations and diagnostics such as wave extraction tools and horizon finders will be required to open up this uncharted and fertile ground of research.  \\vspace{5pt} \\noindent {\\bf Acknowledgements:} U.S. acknowledges support from the Ram{\\'o}n y Cajal Programme and Grant No.~FIS2011-30145-C03-03 of the Ministry of Education and Science of Spain, the FP7-PEOPLE-2011-CIG Grant {\\em CBHEO 293412}, NSF grants PHY-0601459, PHY-0652995 and the Sherman Fairchild Foundation to Caltech. E.B.'s research is supported by NSF grant PHY-0900735 and by NSF CAREER grant PHY-1055103. This work was supported by the DyBHo-256667 ERC Starting Grant, and by FCT-Portugal through projects PTDC/FIS/098025/2008, PTDC/FIS/098032/2008, CTE-AST/098034/2008, CERN/FP/116341/ 2010 and CERN/FP/123593/2011. This research was supported by the DyBHo--256667 ERC Starting Grant, the NRHEP--295189 FP7--PEOPLE--2011--IRSES Grant by allocations through the TeraGrid Advanced Support Program at the San Diego Supercomputing Center and the National Institute for Computational Sciences under grant PHY-090003, the Centro de Supercomputaci{\\'o}n de Galicia (CESGA) under project numbers ICTS-CESGA-175, ICTS-CESGA-200 and ICTS-CESGA-221, the Barcelona Supercomputing Center (BSC) under project AECT-2011-2-0006, AECT-2011-3-0007, AECT-2012-1-0008, the DEISA Extreme Computing Initiatives DECI-6 and DECI-7 PRACE-2IP Project DyBHO."
    },
    "1107/1107.1392_arXiv.txt": {
        "abstract": "We present our near infrared ESO-NTT $K_S$ band observations$^1$ of the field of \\igr\\ which show no extended emission consistent with the SNR but in which we identify a new counterpart,  also visible in \\spitzer\\ images up to 24\\,$\\mu$m,  at the position of the X-ray point source, \\cxopt. Multi-wavelength spectral modelling shows that the data are consistent with a reddened and absorbed single power law over five orders of magnitude in frequency. This implies non-thermal, possibly synchrotron emission that renders the previous identification of this source as a possible pulsar, and its association to the SNR, unlikely;  we instead propose that the emission may be due to a blazar viewed through the plane of the Galaxy. ",
        "introduction": "\\label{section:introduction}    \\let\\thefootnote\\relax\\footnotetext{\\noindent$^1$ Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere, Chile under ESO program 084.D-0535 (P.I. Chaty)}    High energy (X-ray, $\\gamma$-ray) emission is observed from numerous astronomical sources and is produced by various processes. Common sources of high energy emission are Galactic sources such as black hole or neutron star systems (e.g., X-ray binaries, pulsars) and supernovae or supernova remnants, and extragalactic sources such as active galactic nuclei (AGN) and gamma-ray bursts (GRBs). At lower energies ($\\lesssim 20$\\,keV), thermal emission may make a significant contribution to the observed flux but at greater energies, such as those observed by the \\emph{International Gamma-Ray Astrophysics Satellite} (\\integral;  15\\,keV - 10\\,MeV; \\citealt{Winkler2003:A&A411L}), other emission mechanisms are required. Prominent, though certainly not the only emission mechanisms at these energy ranges are synchrotron and inverse Compton radiation which both produce broad-band spectra visible over many orders of magnitude in frequency.  In this Letter we collate data of one such \\integral\\ detected, high energy object, \\igr,  from various sources including published catalogs, archived images and our own ESO-NTT observations (section \\ref{section:observations}). We identify the multi-wavelength counterparts of the point source and construct a spectral energy distribution (SED; section \\ref{section:spectra}) in an attempt to understand its nature.   \\noindent {\\bf \\igr}\\label{section:introduction-igr}\\\\ \\noindent The X-ray source \\igr\\ was initially discovered by INTEGRAL and published in the Third (III) and subsequently Fourth (IV) IBIS/ISGRI Soft Gamma-Ray Survey Catalog \\citep{Bird2007:ApJS170,Bird2010:ApJS186}, though at slightly different positions. In an attempt to refine the position (see Table\\,\\ref{table:xrays} for this and subsequent X-ray positions), the \\emph{Swift} X-ray telescope (XRT; \\citealt{burrows2005:SSRv120}) observed the field \\citep{Landi2007:ATel.1323}. These authors identify a point source (henceforth \\swft) at the edge of the original \\integral\\ III error circle as well as possible diffuse emission (which, after examination of the XRT image, we note is within that error circle). The position of the point source was further refined by \\cite{Tomsick2009:ApJ701} through  a 4.7\\,ks  \\chandra\\ observation of the field. In addition to detecting the point source, \\cxopt, they confirmed an extended source, $\\sim7$ times brighter, at the \\integral\\ III position (\\cxoex, see figure\\,\\ref{fig:24um}). At the XRT point source position, \\cite{Landi2007:ATel.1323} had also noted a USNO-B1.0 (0574-0773466, $R \\sim 15.5$; \\citealt{Monet2003:AJ125}), 2MASS (J17443749-3232197, $K=9.100 \\pm 0.026$; \\citealt{Skrutskie2006:AJ.131}) object but the sub-arcsecond accuracy of the \\chandra\\ position eliminates it as a possible counterpart; though the X-ray point source is at the edge of the 2MASS point spread function (PSF).   Due to the fact that the more recent \\integral\\ IV position includes both the point source and extended emission within its error circle we shall henceforth only use \\igr\\ to refer to the field and, for clarity, we shall refer to the extended emission as \\cxoex. Given the positional coincidence and point source nature of \\swft\\ and \\cxopt\\ we shall henceforth assume they are one source which we shall refer to by the \\chandra\\ moniker. Based on analysis of the \\chandra\\ spectrum, \\cite{Tomsick2009:ApJ701} proposed the extended emission as a supernova remnant (SNR), though they did not detect evidence of a pulsar wind nebula (PWN). They also suggested a tentative association of the point source with the SNR, hypothesizing that it may be an isolated neutron star that received a kick when the supernova occurred.      \\begin{table} \\centering \\caption{X-ray positions and 90\\% uncertainties for the different X-ray sources in the field.} \\label{table:xrays} \\begin{tabular}{l l l l}% \\hline\\hline Source & RA & Declination & Error \\\\ \\hline \\igr$^{\\rm{III}}$   & 17:44:54.96 & -32:33:00 & 2.2\\arcmin \\\\ \\igr$^{\\rm{IV}}$   & 17:44:47.76 & -32:32:16 & 3.2\\arcmin \\\\ \\swft$^{*}$  & 17:44:37.23 & -32:32:24.3 &  2.4\\arcsec \\\\  \\cxoex$^e$ & 17:44:53 & -32:32:54 &   \\\\ \\cxopt & 17:44:37.34 & -32:32:23.0 &  0.64\\arcsec \\\\ \\pbc   & 17:44:52.6  & -32:31:01   &   4.56\\arcmin \\\\  \\hline \\end{tabular} \\begin{list}{}{} \\item[] $^{\\rm{III/IV}}$ Position from the Third/Fourth IBIS/ISGRI Soft Gamma-Ray Survey Catalog. $^*$ Refined in this paper from  original position of 17:44:37.46 -32:32:20.2 (4\\arcsec\\ error; \\citealt{Landi2007:ATel.1323}). $^e$ Extended source of radius $\\sim$ 3\\arcmin. \\end{list} \\end{table} ",
        "conclusions": "\\label{section:conclusion}  On the basis of positional coincidence and common spectral slope, we have identified a new infrared counterpart to the X-ray point source, \\cxopt\\ in the field of  \\igr, visible from 2.2 to 24\\,$\\mu$m. Multi-wavelength spectral modelling shows that the data are consistent with a reddened and absorbed single power law over five orders of magnitude in frequency. This implies non-thermal, possibly synchrotron emission that renders the previous suggestion that this source may be a pulsar, and  its association to the extended SNR emission, unlikely.  We propose that the emission may be due to a blazar viewed through the plane of the Galaxy, and we suggest a number of tests of this hypothesis."
    },
    "1107/1107.6045_arXiv.txt": {
        "abstract": "Encoded within the morphological structure of galaxies are clues related to their formation and evolutionary history.  Recent advances pertaining to the statistics of galaxy morphology include sophisticated measures of concentration (C), asymmetry (A), and clumpiness (S). In this study, these three parameters (CAS) have been applied to a suite of simulated galaxies and compared with observational results inferred from a sample of nearby galaxies.  The simulations span a range of late-type systems, with masses between $\\sim$10$^{10}$~M$_\\odot$ and $\\sim$10$^{12}$~M$_\\odot$, and employ star formation density thresholds between 0.1~cm$^{-3}$ and 100~cm$^{-3}$. We have found that the simulated galaxies possess comparable concentrations to their observed counterparts. However, the results of the CAS analysis revealed that the simulated galaxies are generally more asymmetric, and that the range of clumpiness values extends beyond the range of those observed. Strong correlations were obtained between the three CAS parameters and colour (B-V), consistent with observed galaxies. Furthermore, the simulated galaxies possess strong links between their CAS parameters and Hubble type, mostly in-line with their observed counterparts. ",
        "introduction": "\\label{sec:intro}  Classification in astronomy has led to many advances which have revealed important information about the Universe in which we live. For example, the classification of stars by their colours and brightnesses in the Hertzsprung-Russel (HR) diagram led to an understanding of stellar structure and evolution.  The HR diagram continues to be employed as a means to to determine the ages and metallicites of stellar populations, whether they are simple stellar populations, as in globular clusters, or in composite populations, such as dwarf galaxies \\citep{Dolphin2003, Monelli2010, Williams2010}.  Even when it is impossible to resolve stars, entire galaxies are classified by their colours and magnitudes. However, galaxies display resolved structure and morphology that stars do not, and using this morphology can reveal much about the evolution of galaxies.  Early methods used for the classification of galaxies were based on morphology.  The \\citet{Hubble1926} Classification System is a method used to categorise galaxies by their morphology, and although it has been the dominant morphological tool since the mid-1920s, over time it has proven to be deficient in several areas. A prominent disadvantage to classifying galaxies based on their morphological features alone is that it is subjective with respect to distance (or resolution) and inclination. Furthermore it has led to the grouping of a wide range of asymmetric galaxies as simply ``irregular''.  As both ground- and space-based imaging has improved, and we probe higher redshifts, it has become apparent that classifying galaxies as ``irregular'' means ignoring a vast amount of morphological information. \\citet{deVaucouleurs1959} attempted to improve upon the \\citet{Hubble1926} Classification System by introducing ``later types'' as a sub-class of galaxies; however, this system still refers to local, axisymmetric galaxies as reference points, and, again, was based solely on morphology.  \\citet{Morgan1957} noted a correlation between the spectra of galaxies and the dominance of their central bulge component.  \\citet{Morgan1958} felt that a shortcoming of the Hubble classification scheme was that it was not a strong indicator of spectral class. In order to rectify this,  \\citet{Morgan1958} devised a galaxy classification system based upon the central light concentration compared to its overall light distribution. This concentration classification scheme involved analysing the flux of galaxies to determine the degree of concentration of the central component.  The purpose of this scheme was to study the colour-concentration correlation in an attempt to identify galactic evolutionary states.  Due to the concentration parameter's direct relation with spectral class and type of stellar population, the concentration parameter was also found to be indicative of formation history and properties, such as velocity dispersion, galaxy size, luminosity, and (more recently) central black hole mass \\citep{Graham2001}.  Extending Morgan's pioneering work, \\citet{Conselice2000b} proposed a new statistical measure of asymmetry, in conjunction with a revised parametrisation of Morgan's concentration index.  This classification system was based on the desire to classify galaxies quantitatively at a range of redshifts.  Asymmetry was defined by rotating a galaxy 180$\\degrees$, about a central axis and finding the absolute sum of the normalised residuals.  This method measures the high frequency structure within each galaxy, whilst also considering the symmetry. The measure of asymmetry focuses primarily on morphological shape, and has a strong correlation with colour, congruent with Morgan's concentration parameter. As such, this property has a correlation with both star formation and merger history.  \\citet{Conselice2003} introduced one additional morphological measure sensitive to high spatial frequency clumpiness. This parameter is defined by comparing a galaxy to a smoothed image of itself.  In doing so, the clumpiness parameter quantifies galaxy morphology in a manner which reflects star forming regions and evolutionary history.  This parameter distinguishes between varying types of galaxies and thus, partially, resembles Hubble's classification system.  However, the clumpiness parameter quantifies galaxy morphology with respect to intrinsic characteristics and not morphology alone.  \\citet{Conselice2003} combined the high spatial frequency clumpiness parameter with both the asymmetry index and the concentration index to complete what is now called the CAS classification system. Underlying the CAS system's empirical nature lies the fundamental physics which governs its individual C, A, and S, components.  The CAS system highlights the intrinsic nature of galaxies by considering their light distributions and spectral class.  The observable properties of galaxies are a function of their merger histories, masses and environments, each of which are reflected in the CAS system. It is fitting to apply this new classification method to the theoretically predicted properties of galaxies and use it as a gauge to determine where our current understanding of galaxy formation and evolution is correct and where it is deficient. Galaxy simulations are an integral tool that can help interpolate between known stages of formation history, and thus further our understanding of galaxy evolution.  Within the observational community, the CAS system has become one of the most widely-used measures of galaxy morphology \\citep{Hern2008, Bertone2009}. Since the completion of the combined CAS system, it has been applied to numerous galaxies including the 113 galaxies observed by \\citet{Frei1996}. However, it has not yet been applied systematically to high-resolution computational galaxy simulations as a method of comparison between observations and simulations. In this work, we develop a software package patterned on the \\citet{Conselice2000b} and \\citet{Conselice2003} CAS system, calibrated on an optimised training set, and applied to high resolution galaxy simulations. We show that the simulated galaxies exhibit both similarities and differences to real galaxies, and highlight the intrinsic physical process that are responsible for the differences.  In \\S\\ref{SimSamp} we describe our simulated galaxy samples. In \\S\\ref{CAS} we discuss the three structural parameters of the CAS system, Concentration (C), Asymmetry (A)and Clumpiness (S). We then compare the simulations with the empirical \\citet{Frei1996} sample of galaxies in \\S\\ref{Comparison}. The results are discussed in \\S\\ref{Discuss} and conclusions drawn in \\S\\ref{Concl}. ",
        "conclusions": "\\label{Concl}  We have applied measures of concentration (C), asymmetry (A) and clumpiness (S) to a sample of simulated galaxies which have a range of masses and morphologies, as well as to a sample of observed galaxies. We have explored the correlations between these parameters as well as between each of them with B-V colour.  In general, reasonable agreement was found between the simulated and observed populations in these relations, although some differences become apparent, as summarised below.  Overall, the trend generated by the observed galaxy sample, with respect to colour (B-V), was approximately replicated by the simulated galaxies, although there is a distinct lack of redder simulated galaxies, which is related to a well-established inability of these simulations to halt late time star formation in early type galaxies \\citep{Kawata2005}.  It was established that Hubble type is strongly linked with the three CAS parameters for the both the observed and simulated galaxy samples. This is expected as later-type galaxies, such as Hubble type Sc-Sd, have star forming regions, resulting in significant internal spatial structure, which manifests itself in significant asymmetries and clumpiness, and possess low concentration values. Conversely, earlier-type galaxies are smoother, more symmetric, and more concentrated, due to their relative lack of star forming regions, structure, and their evolutionary histories.  The concentration parameter was found to be robust and easily replicated. The values obtained for the simulated galaxies demonstrated an approximately equal range of values compared with their observed counterparts. The fewer number of simulated galaxies with high concentration merely reflected the relative paucity of early-type galaxies compared with the observed sample.  It was found that the range of asymmetries of the simulated galaxies did not extend to as low values as those of the most symmetric observed galaxies. This may be related to the ongoing star formation in the simulated galaxies, the same process which results in a dearth of red galaxies in the simulations. It is generally the early-type observed galaxies which are highly symmetric. Nevertheless, it remains a challenge for simulators to create highly-symmetric galaxies within a hierarchical structure formation paradigm.  A single dwarf galaxy simulation also shows excessive asymmetry compared to the observed sample. This is due to the existence of a large star forming complex offset from the centre of that simulation. Detailed exploration of the resolution of star forming complexes and star clusters within simulations is beyond the scope of this study, but it appears that the physical size of the star forming complex in at least one of the simulations is larger than observed, at least within the Frei a et al. (1996) sample.  As it is a well known issue that simulated galaxies have disproportionately large bulge components \\citep{Stinson2010}, it was decided that a central region, pertaining to the size of the bulge component, would be removed for all galaxies in the calculation of the clumpiness. Following this, we found that the simulated and observed galaxies showed similar spreads in the values of clumpiness. The early-type disc galaxies from the MUGS sample, which have lower rates star formation occurring in the disk at $z=0$, have low S, while the later-type simulated galaxies continue to have high levels of star formation, and high values of S, consistent with their observed counterparts. Analysis of the clumpiness values generated by the simulated galaxies also demonstrated that several simulated galaxies were excessively clumpy when viewed at an inclination angle of 90$\\degrees$. At least in relation to the dwarfs, this is likely related to high frequency power of these galaxies, when analysed at an edge-on inclination, a characteristic that is not consistent with observations (cf., Pilkington et~al. 2011, for the high-frequency analysis of the ISM of these dwarfs)."
    },
    "1107/1107.3529.txt": {
        "abstract": "Light dark matter (DM) with a large DM-nucleon spin-independent cross section and furthermore proper isospin violation (ISV) $f_n/f_p\\approx-0.7$ may provide a way to understand the confusing DM direct detection results. Combing with the stringent astrophysical and collider constraints, we systematically investigate the origin of ISV first via general operator analyses %from the DM and up/down quark couplings. and further via specifying three kinds of (single) mediators: A light $Z'$ from chiral $U(1)_X$, an approximate spectator Higgs doublet (It can explain the $W+jj$ anomaly simultaneously) and color triplets. In addition, although $Z'$ from an exotic $U(1)_X$ mixing with $U(1)_Y$ generating $f_n=0$, we can combine it with the conventional Higgs to achieve proper ISV. As a concrete example, we propose the $U(1)_X$ model where the $U(1)_X$ charged light sneutrino is the inelastic DM, which dominantly annihilates to light dark states such as $Z'$ with sub-GeV mass. This model can address the recent GoGeNT annual modulation consistent with other DM direct detection results and free of exclusions. ",
        "introduction": "The way that the dark matter (DM) interacts with visible matters is a puzzle due to the absence of confirmative experimental results. The recent possible progresses made on direct detections may shed light on it. On the one hand, the DAMA/LIBRA~\\cite{DAMA} and CoGeNT~\\cite{CoGeNT} experiments report events, which point to a light DM (LDM) having mass $\\sim$8 GeV and a rather large spin-independent (SI) recoil cross section with nucleons, $\\sigma_{n}^{\\rm SI}\\sim10^{-4}$ pb (The supscript SI will be implied). On the other hand, the XENON~\\cite{XENON} and CDMS~\\cite{CDMSII} experiments report null results, which challenges the CoGeNT/DAMA results. Moreover, for the CoGeNT/DAMA favored DM there may be some expected but not found signals from astrophysics or colliders. So how to reconcile these results guides us to identify some of the DM properties.  DM-nucleon interactions with isospin violation (ISV) may provide a way to reconcile various direct detection results~\\cite{ISV}. DM ISV is not a novel phenomena~\\cite{ISVO,Kang:2010mh}, and it arises when the DM-proton and DM-neutron interactions have different strengths, namely $f_{p}\\neq f_n$. In particular, if $f_n/f_p<0$, the DM-proton and DM-neutron scattering amplitudes will destructively interfere, leading to a cancellation in the DM-nucleus scattering amplitude. The degree of cancelation varies with the target nucleus used in different experiments. Ref~\\cite{ISV} showed that if $f_n/f_p=-0.7$, we can not only substantially weaken the XENON100 constraint but also make the CoGeNT- and DAMA-region overlap. This scenario has tension with the CDMS-Ge experiment which uses the same nucleus as the CoGeNT experiment. But if we only consider the annual modulation results (The observed CoGeNT annual modulation has significance of 2.8$\\sigma$~\\cite{CoGeNTA}), the inelastic DM (iDM) scenario~\\cite{iDM} is able to enhance the modulation and thus reduce the tension between the CDMS-Ge and CoGeNT experiments. Ref.~\\cite{Frandsen} found that by taking $f_n/f_p=-0.7$ and the quenching factor $Q_{Na}$ close to its upper-limit $0.43$, one can address all the confusing experimental results via an iDM $\\sim$10 GeV with a mass splitting $\\delta\\simeq 15$ keV. Additionally, right now the direct detection experiments may have reached the level to measure the individual strength of $f_n$ and $f_p$~\\cite{Pato:2011de}. In a word, it is worthy of studying the origin of DM ISV systematically.  %Inspired by the aforementioned facts, in this work we make rather thorough In our analyses, we focus on the scalar and fermionic DMs, which are required to present the following features: \\begin{itemize} \\item A proper ISV $f_n/f_p\\approx -0.7$. How to get this ISV is highly non-trivial. We will show that conventional mediators like the Higgs boson, $Z$ boson and squarks fail to accommodate this value, at least for single mediator case. Thus we need to investigate new mediators beyond them. \\item A large DM-nucleon SI scattering cross section $\\sigma_{n}\\sim  10^{-2}$ pb. In the ISV scenario, the nucleus amplitude is reduced by the destructive interference, so $\\sigma_{n}$ is required to about 2 orders larger than the conventional scenario. That large $\\sigma_n$ will bring tensions with some astrophysical or collider constraints. \\item The right DM relic density $\\Omega_{\\rm DM}h^2\\simeq0.11$. We allow DM annihilation channels whose final states are not the standard model (SM) light fermions. Such channels actually appear in the complete models where LDM can annihilate into light dark sector states. This point will be very helpful somewhere. \\item If possible, the DM models should have connection with other new physics. Especially, we try to account for the recent Tavetron CDF W+2jets anomaly~\\cite{Wjj}. \\end{itemize}   The paper is organized as follows. In Section~II, we make some general  analysis on the dark matter with ISV based on effective operators. In Section~III,  we turn to their possible simple UV origin by specifying the mediators. Next, a concrete model with sneutrino iDM is presented. Conclusion and discussions are given in Section~\\ref{CD}, and finally some useful formulas are collected in the Appendices. ",
        "conclusions": "\\label{CD}   The light dark matter models with ISV $f_n/f_p\\approx-0.70$ and a large DM-nucleon spin-independent cross section $\\sigma_n\\sim{\\cal O} (0.01)$ pb may provide a way to understand the confusing direct detection experimental results. Combining with the stringent astrophysical and collider constraints, we can further deduce the DM properties. In this work, we investigated the possible origin of ISV based on effective analysis, and found that ISV must arise from the DM and first-family quarks couplings. To further explore their UV origins, we considered the operators as a result of integrating out the following mediators: \\begin{description} \\item[$Z'$ from $U(1)_X$] The $U(1)_X$ must be chiral, and the light $Z'$ is strongly favored. \\item[Spectator Higgs doublets] Conventional Higgs doublet mediates interactions preserving isospin, so we introduce (approximate) spectator Higgs doublet whose couplings to the SM quarks are free parameters. Such a Higgs doublet can be additionally used to explain Tevatron CDF $W+jj$ anomaly. \\item[Color triplets]  Combining the squarks in the MSSM with exotic $U(1)_X$ (quarks strongly charged under it), we found that the light $\\wt B_X$ LSP can generate proper ISV via the first-family squark mediation. This  scenario is economic furthermore suffers no flavor problem. \\item[Exotic $Z'$ plus Higgs] For a SM-neutral $U(1)_X$ having kinetic mixing with $U(1)_Y$, its   light gauge boson $X_\\mu$ only mediates DM-proton interaction. Combining it with a conventional Higgs mediator, we can obtain the desired ISV. Similar point is also adopted in Ref.~\\cite{DelNobile:2011je} during the completion of this work. \\end{description} As a realistic model building, we propose the MSSM with low scale seesaw mechanism and sub-GeV scale $U(1)_X$ gauge group extension. In this model a light sneutrino plays the role of isospin-violating-iDM to explain the CoGeNT annual modulation, in consistent with other detections results and various bounds.  ."
    },
    "1107/1107.4106_arXiv.txt": {
        "abstract": "The recent distance determination allowed precise estimation of the orbital parameters of Cyg X-1, which contains a massive $14.8\\msun$ BH with a $19.2\\msun$ O star companion. This system appears to be the clearest example of a potential progenitor of a BH-NS system.  We follow the future evolution of Cyg X-1, and show that it will soon encounter a Roche lobe overflow episode, followed shortly by a Type Ib/c supernova and the formation of a NS. It is demonstrated that in majority of cases ($\\gtrsim70\\%$) the supernova and associated natal kick disrupts the binary due to the fact that the orbit expanded significantly in the Roche lobe overflow episode. In the reminder of cases ($\\lesssim30\\%$) the newly formed BH-NS system is too wide to coalesce in the Hubble time. Only sporadically ($\\sim1\\%$) a Cyg X-1 like binary may form a coalescing BH-NS system given a favorable direction and magnitude of the natal kick. If Cyg X-1 like channel (comparable mass BH-O star bright X-ray binary) is the only or dominant way to form BH-NS binaries in the Galaxy we can estimate the empirical BH-NS merger rate in the Galaxy at the level of $\\sim 0.001$ Myr$^{-1}$. This rate is so low that the detection of BH-NS systems in gravitational radiation is highly unlikely, generating Advanced LIGO/VIRGO detection rates at the level of only $\\sim 1$ per century. If BH-NS inspirals are in fact detected, it will indicate that the formation of these systems proceeds via some alternative and yet unobserved channels. ",
        "introduction": "Estimates for the rates of gravitational radiation (GR) sources from coalescing degenerate binaries are typically performed with population synthesis methods (e.g., Lipunov, Postnov \\& Prokhorov 1997; Bethe \\& Brown 1998; De Donder \\& Vanbeveren 1998; Bloom, Sigurdsson \\& Pols 1999; Fryer, Woosley \\& Hartmann 1999; Nelemans, Yungelson \\& Portegies Zwart 2001, Voss \\& Tauris 2003 or more recently Belczynski et al. 2010). These studies attempt to explain the past evolution of the given observed binary or stellar population and put some constraints on the physics of stellar/binary evolution (e.g., Valsecchi et al. 2010). Nevertheless, there are often critical model parameters that are poorly constrained.  Recently, Bulik, Belczynski \\& Prestwich (2011) have taken a different approach, examining specific binary systems with well established parameters and investigating their future evolution. If a binary is chosen close to the end of its life (e.g., the formation of double compact object) such a method has potentially great predictive power as many unknowns relating to its prior binary and stellar evolution can be avoided. In particular, Bulik et al. (2011) considered two high mass X-ray binaries (HMXBs), IC10 X-1 and NGC300 X-1, and showed that these systems will soon form close black hole + black hole (BH-BH) systems that will merge within Hubble time and produce strong GR signature. This provided an empirical lower limit of detection chances for the current GR instruments without direct reference to population synthesis methods.  In this study, we consider the future evolution of one of the most interesting binaries known in our Galaxy: Cyg X-1. Recently, the distance to this system was determined by radio parallax and other methods (Reid et al. 2011; Xiang et al. 2011), allowing the basic parameters of this binary to be  firmly established (Orosz et al. 2011). This HMXB hosts one of the most massive ($15 \\msun$) Galactic BHs in a close orbit around a massive ($20 \\msun$) O star. Since the companion is in the mass range for neutron star formation, we have selected this system to investigate yet another unobserved population of potential GR sources: black holes + neutron star (BH-NS) systems. ",
        "conclusions": "So far we have discussed only one binary, Cyg X-1, as a potential progenitor of BH-NS system. Tomsick \\& Muterspaugh (2010) list several known Galactic HMXBs with a NS and a companion massive enough to potentialy produce a BH: Vela X-1 ($24\\msun$), XTEJ1855-026 ($25\\msun$), 4U1907+09 ($28\\msun$) and GX301 ($40\\msun$). The expansion of massive companions will eventually lead to a RLOF. These systems due to extreme mass ratio will evolve into common envelope that will result in a merger of a NS with its companion aborting the BH-NS formation. Such outcome follows from the fact that the mass of a NS is so small in respect to the mass of the envelope of a companion that there is not enough orbital energy to eject companion envelope (e.g., Webbink 1984). An extragalactic system LMC X-3 hosting a $10\\msun$ BH was belived to have a companion massive enough to form a NS. However, recent spectral analysis that takes into account irradiation of the companion, indicates that it is a B5 dwarf setting its mass at about $5\\msun$ (Val-Baker, Norton \\& Negueruela 2007) and makes it a WD progenitor.  HMXBs like Cyg X-1 are wind fed, and thus their non-degenerate stars do not fill their Roche lobes. It is therefore curious that those binaries for which good system parameters are known tend to be close to RLOF. Cyg X-1 is an example: our analysis implies that the system will begin RLOF in $0.2$ Myr, after a lifetime about 50 times longer than that.  An even more extreme case is that of LMC X-1, in which an $11\\msun $ black hole is found in a 3.9 day orbit with a $32\\msun $ companion.  The companion currently fills over 90\\% of its Roche lobe, implying that the system will undergo RLOF within $\\approx 0.1$ Myr (Orosz et al. 2009). Note that the high mass of the companion implies that RLOF will result in unstable mass transfer and the likely formation of a common envelope system --- this LMC X-1 is unlikely to produce any sort of double degenerate binary. There are only a few HMXBs with precisely known system parameters, so this tendency to be alarmingly near RLOF may be a statistical anomaly. Alternatively, it may be a selection effect: a system in which the companion is far from RLOF will have a smaller fraction of its stellar wind accrete onto the compact object, and thus a lower X-ray luminosity.  Nevertheless, it may be worth entertaining the idea that this effect is neither a statistical glitch nor an observational bias, but that there is some unknown physical process that halts the growth of the companion star near the Roche lobe boundary.  We note that the gravitational potential becomes shallow near the $L_1$ point, and  X-ray irradiation becomes a bigger effect for companion stars that present a relatively large cross section. Both of these effects may dramatically change the surface structure and stellar winds of the companion star as it approaches the Roche lobe, conceivably in a self-limiting way.  Such winds might be observable through the strength, shape and variability of emission lines associated with the wind.  If such an effect is present, the evolutionary path of the binary system may prove to be quite different from what is commonly assumed, and what we have assumed above.  To evaluate the possible effect this might have on rates of GR sources, we consider how the future evolution of an HMXB like Cyg X-1 might proceed if it never reaches RLOF.  As a limiting case, we constrain the orbital period to remain at its currently observed value of $P_{\\rm orb}=5.6$d. In fact, we expect the orbital period to increase somewhat (even without RLOF) as mass is lost from the system in stellar wind. In this case, the survival through the supernova and the formation of the close BH-NS system is expected in $8.4\\%$ of cases and corresponds to Advanced LIGO/VIRGO detection rate of $\\sim 1$ per decade. So despite the fact that we have violated (in favor of producing close BH-NS systems) our current understanding of the stellar evolution we still do not get enough BH-NS mergers to expect detection in gravitational waves.  Thus we find that if indeed BH-NS mergers are observed as GR sources, their immediate precursors will not be systems like Cyg X-1 that are currently observed."
    },
    "1107/1107.4330_arXiv.txt": {
        "abstract": "Microwave Kinetic Inductance Detectors (MKIDs) have shown great potential for sub-mm instrumentation because of the high scalability of the technology. Here we demonstrate for the first time in the sub-mm band (0.1\\ldots2~mm) a photon noise limited performance of a small antenna coupled MKID detector array and we describe the relation between photon noise and MKID intrinsic generation-recombination noise. Additionally we use the observed photon noise to measure the optical efficiency of detectors to be $0.8\\pm 0.2$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1107/1107.0725_arXiv.txt": {
        "abstract": "The numerous and massive young star clusters in blue compact galaxies (BCGs) are used to investigate the properties of their hosts. We test whether BCGs follow claimed relations between cluster populations and their hosts, such as the the fraction of the total luminosity contributed by the clusters as function of the mean star formation rate density; the $V$ band luminosity of the brightest youngest cluster as related to the mean host star formation rate; and the cluster formation efficiency (i.e., the fraction of star formation happening in star clusters) versus the density of the SFR. We find that BCGs follow the trends, supporting a scenario where cluster formation and environmental properties of the host are correlated. They occupy, in all the diagrams, the regions of higher SFRs, as expected by the extreme nature of the starbursts operating in these systems. We find that the star clusters  contribute almost to the 20 \\% of the UV luminosity of the hosts. We suggest that the BCG starburst environment has most likely  favoured the compression and collapse of the giant molecular clouds, enhancing the local star formation efficiency, so that massive clusters have been formed. The estimated cluster formation efficiency supports this scenario. BCGs have a cluster formation efficiency comparable to luminous IR galaxies and spiral starburst nuclei (the averaged value is $\\sim 35$ \\%) which is much higher than the  8-10 \\% reported for quiescent spirals and dwarf star-forming galaxies. ",
        "introduction": "Luminous blue compact galaxies (BCGs) are star-forming systems with high specific star formation rate \\citep{2001A&A...374..800O, 2005NewA...11..103S}. It has been suggested that BCGs may have accounted for $\\sim$ 40\\% of the total star formation rate (SFR) density at redshift $0.4<$z$<1.0$ playing an important role in the star formation history (SFH) of the Universe \\citep{1997ApJ...489..559G}.  However, in the local Universe, their contribution has dropped drastically and  luminous BCGs with high SFR have become rare objects \\citep{1997ApJ...489..559G, 2004ApJ...617.1004W}. At a distance of 100 Mpc (z$\\sim$0.03), we count only a handful number of BCGs with a SFR higher or close to 5 $\\msun$yr$^{-1}$.  Among the systems included in this work, Haro 11 \\citep{A2010}, ESO 185-IG13 \\citep{A2010c}, and Mrk\\,930 \\citep{A2011a} are starburst BCGs with SFRs exceptionally high for local irregular galaxies. Local luminous BCGs display physical conditions (low metallicity and dust content), morphologies (compactness of the starburst regions), and feedback mechanisms \\citep[e.g.][]{2009AJ....138..923O} similar to their high redshift counterparts.  Numerical simulations and theoretical arguments based on the Lambda Cold Dark Matter model of the Universe predict that smaller galaxies formed first and were then accreted into more massive systems  \\citep[hierarchical growth,][]{2000MNRAS.319..168C, 2005Natur.435..629S, 2005ApJ...631..101P}. Therefore, it is expected that the primordial galaxies were chemically unevolved dwarf galaxies, the so-called \"building-blocks\" systems, which accrete and merge into more massive units. Recent observations of galaxies at very high redshifts  (z$\\sim7$, i.e. at the reionization epoch) have revealed  compact single or double nuclei systems with extended nebular features \\citep{2010ApJ...709L..21O}. Some of these z$\\sim7$ objects have also  been detected at longer IR wavelengths. These data have been used to constrain and study the spectral energy distributions (SEDs) of these high-z systems. The inferred SFRs are between 5 to 20 $\\msun$yr$^{-1}$ \\citep{2010ApJ...713..115G}. In general, they have estimated stellar masses of $10^{8-9}$ $\\msun$ and fainter UV luminosities than lower redshift Lyman break galaxies \\citep{2010ApJ...719.1250F}, i.e. their properties resemble those of the luminous BCGs.  The studies of BCGs can, therefore, give important insights in understanding how star formation proceeds in dwarf starburst galaxies and possibly in high-redshift systems, with spatial and spectral resolution impossible to achieve for the latter with the current facilities. High-resolution imaging data of BCGs revealed that the starburst regions in these galaxies are formed by massive and young star cluster complexes \\citep[e.g.][]{2003A&A...408..887O, A2010, A2010c, A2011a}. The peak age of the star cluster distributions and the estimated ages of the starbursts in these systems are in good agreement. The host morphologies suggest that the galaxies have recently undergone a merger or interaction event, which has likely refurbished the galaxy with  metal-poor gas, and triggered a vigorous starburst episode.  Star clusters are a natural outcome of the star formation process (see \\citealp{2003ARA&A..41...57L} for a review). However, for very young stellar systems (e.g. a few Myr or less), it is not trivial to make a clear definition of a cluster \\citep{2010MNRAS.409L..54B}. Usually, it is assumed that stars form in a clustered fashion and that, after roughly 10 Myr, 90 \\% of clusters are destroyed due to  gas expulsion.  \\citet{2010MNRAS.409L..54B}  suggest that all the stars form in a continuos hierarchy, even the dense regions and estimate that, in the solar neighbourhood,  only $\\sim 26 \\%$ of young stellar objects (new born stars) are located in denser regions, i.e. embedded star clusters. The remaining stars form in associations and agglomerates following a hierarchical continuum distribution with clusters at the bottom of the process. \\citet{2010ARA&A..48..431P} and \\citet{2010MNRAS.tmpL.168G} defined an empirical relation to separate clusters from loose associations. For the latter the crossing time is much larger than the age of the stars, i.e. they are dynamically unbound systems. They observed that the division between bound and unbound systems becomes more clear once the gas-expulsion phase is over ($\\sim$ 10-20 Myr).  Cluster formation at high-redshift is almost unknown. In the local Universe, we observe the evolved counterparts, i.e., the globular clusters (GCs). However, there is no direct evidence that the young star clusters we observe locally will eventually evolve into GCs \\citep{2007ChJAA...7..155D}.  If one looks at the number of GCs per mass (luminosity) bins (i.e., $dN(M)/dM=CM^{\\eta}$) the distribution can be fitted by two power-laws with a slope of $\\sim-2.0$ at the massive (luminous) ends, a break corresponding to the $M_c\\sim2\\times10^5 \\msun$, and a flattening ($\\eta\\sim-0.2$ which variates from system to system) at lower masses. Several possible scenarios are addressed by theoretical studies to explain the origin of the GC mass (luminosity) function: dynamical evolution, i.e. due to preferential disruption of the low mass systems; or  a primordial origin, i.e. the GC mass function has been established at the time when the GCs formed \\citep{2007ChJAA...7..155D}. Cosmological simulations seem to suggest that GCs may have formed in dark matter halos \\citep[e.g.][]{2005ApJ...623..650K, 2005ApJ...619..258M}. If this is the case, there is no connection between the YSCs forming at the present time and the ancient GCs. However, some recent works \\citep[see][for a review on GCs]{2006ARA&A..44..193B} suggest that cluster formation in dwarf galaxies may be the key to explain observed properties of the ancient GCs (blue versus red, i.e. metal poor versus metal rich). \\citet{2010ApJ...718.1266M} recovered a bimodal metallicity distribution as a product of cluster formation in different phases of galaxy evolution. Cluster formation in dwarf galaxies produced the blue, metal-poor GCs.  Dwarf systems were successively accreted to form more massive systems. During the merger and formation of these massive galaxies the more metal-rich clusters were formed. Accretion of dwarf galaxies together with their GC systems is also one of the proposed scenarios by \\citet{2011A&A...525A..20C} to explain why the younger GCs in S0 type of galaxies appear blue instead of the expected metal-rich populations \\citep{2006ARA&A..44..193B}.  In the present work, we will investigate how cluster formation has proceeded in BCGs. We will test whether BCGs follow the cluster-host relations available in the literature and constrained using local star-forming galaxies, like spirals and dwarfs. These results will be used to constrain the environmental properties of BCGs and whether star formation operates on similar modes even under extreme conditions.  The paper is organized as follow: In Section 2, we present the 5 BCG targets used in this work. In Section 3, we first discuss the uncertainties which affect the analysis. In the second part of this section, we show the three cluster-host relations including the BCG sample. A discussion of the results is presented in Section 4. Here, we also discuss possible similarities between BCGs and high redshift galaxies. Conclusions are summarized in the last section.   \\section[]{The data} \\label{data-sample}  In this section we shortly introduce the BCGs included in the analysis. Some of these BCGs have been studied in a series of 3 papers: Haro 11 analysis is presented in \\citet[][]{A2010}; ESO 185-IG13 (ESO 185) in \\citet[][]{A2010c}; and Mrk\\,930 in \\citet[][]{A2011a}. We refer to those papers for details on the analysis of the data used in this work.  \\subsection{ESO 338-IG04}  The analysis of the star cluster population in ESO 338-IG04 (ESO 338) has been presented in \\citet{2003A&A...408..887O}. The masses have been obtained from models that assume a Salpeter initial mass function \\citep[IMF,][]{1955ApJ...121..161S}. We show, in Figure~\\ref{CMF_eso338}, the cluster formation history during the last 40 Myr of galaxy evolution. Using the age and mass estimates from \\citet{2003A&A...408..887O}, we assumed that the analysis is complete in detecting clusters more massive than $5\\times10^3$ $\\msun$ formed during the last 40 Myr. A power law cluster mass function with index $-2.0$ has then been used to extrapolate the total fraction of mass in clusters including objects with masses between $10^2\\leq$M$\\leq5\\times10^3 \\msun$. Following \\citet[][]{2004MNRAS.350.1503W}, we assume for simplicity that a cluster population forms every 10 Myr and estimate the cluster formation rate (CFR) in the galaxy. In agreement with \\citet{2003A&A...408..887O}, we observe a cluster formation enhancement between 20 and 30 Myr ago. However, the peak of cluster formation is younger than 10 Myr, similarly to what is found in Haro 11, ESO 185, and Mrk\\,930. ESO 338 is forming roughly 1.6 $\\msun$yr$^{-1}$ of stars in clusters.  We defined the cluster formation efficiency (CFE, indicated hereafter also as $\\Gamma$) as that the fraction of stars formed in star clusters. If we compare the CFR to the mean SFR happening in the galaxy, we obtain the CFE$=$CFR/SFR \\citep[see G10 and][hereafter B08]{2008MNRAS.390..759B}. Using a SFR of 3.2 $\\msun$yr$^{-1}$ \\citep{2001A&A...374..800O}, we find that $\\Gamma \\sim$50 $\\pm$ 10 \\%, i.e. 50 \\% of  the ESO 338 stars are formed in clusters.  \\begin{figure} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{CMF_40myr_ESO338.ps}}} \\caption{Cluster formation rates during the last 40 Myr of starburst activity in ESO 338. The filled dots connected by the green dotted line show the observed CFR derived from clusters more massive than $10^4 \\msun$. The squares show the derived CFR, if the total mass in clusters less massive than $10^4 \\msun$ is extrapolated using a CMF with index $-2.0$ down to $10^2 \\msun$. The data used are from \\citet{2003A&A...408..887O}.} \\label{CMF_eso338} \\end{figure}   \\subsection{SBS 0335-052E}  The nuclear region of this extreme metal-poor galaxy is dominated by 6 massive young star clusters which have been referred to in the literature as super star clusters (SSCs). These SSCs have been extensively studied before (e.g., \\citealp{2008AJ....136.1415R}; \\citealp{2009ApJ...691.1068T}; \\citealp{A2010b} among the most recent published works on this subject).   \\begin{figure} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{U-B_vs_B-V.ps}}} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{mas_vs_age0004_f0.01_n4.ps}}} \\caption{In the top panel the color-color diagram of the whole cluster population detected in SBS 0335-052E is shown. Blue dots  indicate the six super massive star clusters, which dominate the starburst in the galaxy and have already been studied in several works (e.g. \\citealp{A2010b} for references). The black dots represent the newly studied lower mass cluster population. The evolutionary track has a metallicity of $Z = 0.0004$ \\citet[see][for details]{A2010b}. Some important ages are labelled. The extinction vector indicates in which direction the clusters move if a de-reddening of  A$_V=1$ mag is applied to the observed colors. A mean photometric error-bar is shown. In the bottom panel, we show the mass-age diagram of the star cluster population. The several lines indicate the detection limits reached in our observed data (shown in the inset).  See main text for details.} \\label{mas-age} \\end{figure}  In the present work we explore the lower mass cluster population of the galaxy. Part of this underlying cluster population has already been revealed by \\citet{1998A&A...338...43P}.  We have access to $FUV$ ({\\it SBC}/F140LP from GO 9470, PI: Kunth) and optical images ({\\it ACS} F220W, F330W, F435W, F550M from GO 10575, PI: Ostlin) for the galaxy. The reduction of the science frames is described in \\citet{2009AJ....138..923O}. The extraction of the cluster candidates has been done using the {\\tt PyRAF} package {\\tt DAOFIND} on the $B$ band frame (F435W). The catalogue has been cleaned by eye of all the detections which didn't show a clear visual counterpart.  With this first catalogue, we have done photometry on all the frames from the $FUV$ to the $V$ (F550M) bands. Applying the same method \\citep[see][]{A2010, A2010c, A2011a} we have included in the final catalogue only cluster candidates with detection in at least 3 filters and a photometric error $\\sigma_m < 0.2$ mag. The photometric properties of the final cluster population are shown in the color-color diagram in Figure~\\ref{mas-age}. The filter combination corresponds to a $U-B$ (F330W$-$F435W) versus a $B-V$ (F435W$-$F550M) color. The colors are not dereddened  but clearly suggest that the cluster candidates are not older than $\\sim 100$ Myr.  The spectral energy distribution (SED) fitting procedure is  described in \\cite{A2010}. The used models are presented in \\citet{A2010b}. The output age and mass of the clusters are shown in the mass-age diagram in Figure~\\ref{mas-age}. In blue dots, we show the 6 SSCs previously analysed. The underlying black dots represent the low mass cluster population that is present in the galaxy, detectable at the detection limits imposed by the data. The masses are smaller than $\\sim 5\\times 10^4$  $\\msun$ and the ages are younger than 50 Myr. It is still under debate whether an old stellar population (older than 100 Myr) exists in SBS 0335-052E \\citep{2001A&A...371..429O} or if this galaxy has recently formed \\citep{1998A&A...338...43P}. In previous star cluster analyses of other BCGs (ESO 338, Haro 11, ESO 185, Mrk\\,930), we have found a trace of some old GCs, supporting the evidence of an old underlying stellar population in those galaxies.  In the case of  SBS 0335-052E, however, the non detection of massive GCs cannot prove/disprove either of the two proposed scenarios. We observe that, because of the detection limits (see inset in Figure~\\ref{mas-age}), our analysis is limited to GCs with masses higher than $5\\times 10^4$ $\\msun$ between 100 Myr and 1 Gyr and even more massive at older ages. Therefore, we cannot exclude that this galaxy has low-mass GCs.  Using the same method as in the case of ESO 338, we infer  the mass in clusters formed during the last 10 Myr in SBS 0335. We observe that the total mass contained in detected clusters younger than 10 Myr is  $1.95\\times 10^6$ $\\msun$. The mass contained in the 5 SSCs (one of them is much older, e.g. from H$\\alpha$ equivalent width, the age is constrained  to $\\sim13$ Myr is constrained, see \\citealp{A2010b}) is roughly 73 \\% of this total mass ($1.42\\times 10^6$ $\\msun$). Assuming a power law cluster mass function and that we are complete in detecting clusters more massive than $5\\times10^3 \\msun$, we estimate a total mass in clusters younger  than 10 Myr of $6.4\\times10^6 \\msun$. The observed CFR in systems more massive than M$>5\\times10^3$ $\\msun$ is 0.2 $\\msun$yr$^{-1}$, while the extrapolated CFR (M$>10^2$ $\\msun$)$=0.64$ $\\msun$yr$^{-1}$. The $\\Gamma$ value of SBS 0335 is  49 $\\pm 12 \\%$, using a SFR of 1.3 $\\msun$yr$^{-1}$. These values have been estimated after that a correction for Salpeter IMF has been applied (see Section~\\ref{sfr_unc}). ",
        "conclusions": "To understand the role of BCGs in the galaxy formation scenario, we have tried to constrain how the star and cluster formation is operating in these systems.  We have looked into the efficiency of the formation of massive star clusters. It is known that  some properties of the cluster population and the mean star formation in the host system are correlated, suggesting that the formation of a cluster is not a local event but intrinsically connected to the mean properties of the galaxy. In general, we find that BCGs follow fairly well these relations, even if their SFRs and cluster properties are more extreme.  We discuss possible uncertainties affecting our results. The relation which is least affected is the M$_V^{\\textnormal{brightest}}$-SFR one, which suggests that the cluster formation efficiency  is higher in BCGs than in quiescent spiral and dwarf starburst galaxies. The same evidence is also suggested by the other two relations, $\\Gamma$-$\\log(\\Sigma_\\mathrm{SFR})$ and T$_L$(U)-$\\Sigma_\\mathrm{SFR}$, which appear to be consistent with the general picture despite the uncertainties of the data.  In particular, we observe that the inclusion of the BCG sample separates the $\\Gamma$-$\\log(\\Sigma_\\mathrm{SFR})$ plane into two regions. Local spiral galaxies and dwarf starbursts as the Magellanic Clouds occupy the area around a mean CFE of  $8 \\pm 5$ \\% in agreement with the prediction made by B08. The BCGs, together with a LIRG and a nuclear starburst region (included in the G10 sample), are in a region with a mean CFE of $35.6 \\pm 10$ \\%. This indicates that the merger event has enhanced the cluster formation in these systems.  We observe that the fraction of light produced by  star clusters in BCGs increases at shorter wavelengths and contribute significantly to the UV and U luminosities. This suggests that clustered regions contribute a substantial fraction to the rest-frame UV light of z$\\sim$1-3 galaxies with  metallicities similar to the BCGs. In studies of Lyman break galaxy analogs at redshift between $\\sim$0.1 and 0.2, \\citet{2008ApJ...677...37O} observed super starburst compact regions which dominate the UV light of these targets.  In our analysis, we find some evidence that those bright clumps are probably unresolved star cluster knots, similarly to the ones observed in BCGs."
    },
    "1107/1107.2516_arXiv.txt": {
        "abstract": "The identification, and subsequent discovery, of fast radio transients through blind-search surveys requires a large amount of processing power, in worst cases scaling as $\\mathcal{O}(N^3)$. For this reason, survey data are generally processed offline, using high-performance computing architectures or hardware-based designs. In recent years, graphics processing units have been extensively used for numerical analysis and scientific simulations, especially after the introduction of new high-level application programming interfaces. Here we show how GPUs can be used for fast transient discovery in real-time. We present a solution to the problem of de-dispersion, providing performance comparisons with a typical computing machine and traditional pulsar processing software. We describe the architecture of a real-time, GPU-based transient search machine. In terms of performance, our GPU solution provides a speed-up factor of between 50 and 200, depending on the parameters of the search. ",
        "introduction": "\\label{introductionSection}  One of the most ambitious experiments of next-generation radio-telescopes, such as the newly inaugurated Low Frequency Array (LOFAR) and the future Square Kilometre Array (SKA) is to explore the nature of the dynamic radio sky at timescales ranging from nanoseconds to years. Surveys for very fast radio transients necessarily generate huge amounts of data, making storage and off-line processing an unattractive solution. On the other hand, real-time processing offers the possibility to react as fast as possible and conduct follow-up observations across the electromagnetic spectrum and even, for some events, with gravitational wave detectors. Here we will consider the case where time-series data (tied-array, or beam-formed data for array telescopes mentioned above) are processed to extract astrophysical radio bursts of short duration. The multitude of potential astrophysical events that may produce such transients, ranging from individual pulses from neutron stars to Lorimer bursts \\citep{lorimer2007}, are discussed in \\cite{cm04}.  In this paper we will mostly be concerned with de-dispersion. This refers to a family of techniques employed to reverse the frequency-dependent refractive effect of the interstellar medium (ISM) on the radio signals passing through it. We also use interstellar scattering to constrain the parameter space of a given search - see \\citep[chapter 4]{LorimerKramer2005} for more details on these effects. From pulsar studies, it is well determined that propagation through the ISM obeys the cold plasma dispersion law, where the time-delay between two frequencies $f_1$ and $f_2$ is given by the quadratic relation \\begin{equation} \\label{dispRelationshipEquation} \\Delta t \\simeq 4.15 \\times 10^6 \\mbox{ ms} \\times \\left( f^{-2}_1 - f^{-2}_2 \\right) \\times \\mbox{DM} \\end{equation} where DM is the dispersion measure in $\\text{pc cm}^{-3}$ and $f_1$ and $f_2$ are in MHz. Astrophysical objects have a particular DM value associated with them, which depends on the total amount of free electrons along the line of sight and, therefore, on the distance to the object from Earth.  Removing the effect of dispersion from astrophysical data is typically done in two ways, depending on the type of data and the requirements of the experiment. For total power filterbank data, where the spectral bandwidth of the observation is typically split into a number of narrow frequency channels, de-dispersion consists of a relative shift in time of all frequency channels according to equation \\ref{dispRelationshipEquation}. This is known as incoherent de-dispersion as it is performed on incoherent data. For baseband data, de-dispersion can be done by convolution of the observed voltage data with the inverse of the transfer function of the ISM. This is known as coherent de-dispersion; it is more accurate in terms of recovering the intrinsic shape of the astrophysical signal but much more demanding in computational power. In this paper, we deal with incoherent de-dispersion, applied to total power data.  De-dispersion is an expensive operation, scaling as $\\mathcal{O}\\left(N^2\\right)$ for brute-force algorithms, with more optimized techniques diminishing this to $\\mathcal{O}\\left(Nlog(N)\\right)$. For any blind, fast transient search, de-dispersion needs to be performed over a range of DM values, practically multiplying the number of operations required by the number of DM values in the search. Typically, thousands of finely spaced DM values need to be considered, so the entire operation becomes extremely compute-intensive.  In the past, searches for fast transients, such as single-pulse pulsar searches, have been performed offline on archival data, using standard processing tools on large computing clusters or supercomputers. The power of such searches is highlighted by the discovery of Rotating Radio Transients (RRATs) \\citep{mll+06}.  Here we discuss how such searches can be performed online by using standard off-the-shelf general purpose graphics processing units (GPUs).  With the introduction of new high level application programming interfaces (APIs) in recent years, such as CUDA \\citep{cudaOnline}, these devices can easily be used for offloading data-parallel, processing-intensive algorithms from the CPU.  The idea of performing transient-related signal processing on GPUs is not a new one. For example, Allal et al. (2009) have created a fully working, real-time, GPU-based coherent de-dispersion setup at the Nan\\c{c}ay Radio Observatory \\nocite{Allal2009}. The digital signal processing library for pulsar astronomy written by van Straten and Bailes (2010)\\nocite{VanStraten2010} can also use GPUs for performing coherent de-dispersion, folding and detection.  \\nocite{Barsdell2010}Barsdell, Barnes \\& Fluke (2010) analyse various foundation algorithms which are used in astronomy and try and determine whether they can be implemented in a massively-parallel computing environment. One of the algorithms they analyse is the brute-force de-dispersion algorithm, and they state that this would likely perform to high efficiency in such an environment, whilst stating that optimal arithmetic intensity is unlikely to be achieved without a detailed analysis of the algorithm's memory access patterns. The incoherent de-dispersion algorithm has indeed been implemented using CUDA as a test-case for the ASKAP CRAFT project (see Macquart et al. and Dodson et al. 2010)\\nocite{Macquart2010,Dodson2010}. ",
        "conclusions": "\\label{conclusionSection}  We have implemented two de-dispersion algorithms, brute-force and subband de-dispersion, using CUDA, which enables data-parallel processing to be offloaded onto any number of connected CUDA-enabled GPUs. This has led to a performance speedup of about two orders of magnitude, between 50 and 200 for certain parameter configurations, when compared to a single-threaded CPU implementation. Detailed comparison with two traditional pulsar processing suites, PRESTO and SIGPROC, confirm our results. Finally, a prototype for a real-time dispersion search pipeline was designed, which reads in a UDP stream of telescope data and performs FFT channelisation and de-dispersion.  Work is ongoing in this project, with plans to add several additional processing modules in the pipeline. Coherent de-dispersion is useful for studying pulsars whose DM value is already known. Other schemes for de-dispersion are also being considered, such as performing chirp analysis in the frequency domain to detect chirps representing dispersed signals. GPUs provide us with enough processing power (per unit cost) to be able to apply processing-intensive algorithms which would otherwise be unfeasible on a conventional CPU system. As a result, we are in a position to carry out real-time searches for dispersed fast-transients with appropriate telescopes at low cost."
    },
    "1107/1107.5100_arXiv.txt": {
        "abstract": "We present the results of a  long-term observation campaign of the extragalactic wind-accreting black-hole X-ray binary LMC~X--1, using the Proportional Counter Array on the Rossi X-Ray Timing Explorer (RXTE). The observations show that LMC~X--1's accretion disk exhibits an anomalous temperature-luminosity relation. We use deep archival RXTE observations to show that large movements across the temperature-luminosity space occupied by the system can take place on time scales as short as half an hour. These changes cannot be adequately explained by perturbations that propagate from the outer disk on a viscous timescale. We propose instead that the apparent disk variations reflect rapid fluctuations within the Compton up-scattering coronal material, which occults the inner parts of the disk. The expected relationship between the observed disk luminosity and apparent disk temperature derived from the variable occultation model is quantitatively shown to be in good agreement with the observations. Two other observations support this picture: an inverse correlation between the flux in the power-law spectral component and the fitted inner disk temperature, and a near-constant total photon flux, suggesting that the inner disk is not ejected when a lower temperature is observed. ",
        "introduction": "LMC~X--1 is (with \\cyg) one of only two known persistently luminous x-ray binaries consisting of a black hole accreting the wind of a massive blue star. The $10.91 \\pm 1.41 M_{\\sun}$ black hole is in a $3.90917 \\pm 0.00005$ day orbit \\citep{lmcx1-orosz-2009} about an O7~III companion \\citep{lmcx1-cowley-1995}. As the companion both drives a strong wind and is far from filling its Roche lobe, wind accretion feeds the black hole \\citep{lmcx1-and-lmcx3-nowak-2000}.  The system, which accretes at an average of $0.16 L_{Edd}$ \\citep{lmcx1-gou-2009} has never been observed in the low/hard state \\citep{lmcx1-and-lmcx3-wilms-2001}.  \\citet{lmcx1-gou-2009} used x-ray data, including some of the {\\it RXTE} observations we use here, to derive a spin parameter $\\alpha = 0.92 (+0.05, -0.07)$ for the black hole in LMC~X--1 using fits to the disk blackbody component of the spectrum.   \\xone's persistent occupation of the high/soft state contrasts with the behavior of \\cyg, which it closely resembles in other respects. Cyg~X--1 harbors a $\\sim 10M_{\\odot}$ black hole \\citep{herrero-1995} in a $5.6$ day orbit \\citep{cygx1-paczynski-1974} about an O9.7 Iab companion \\citep{cygx1-bolton-1972}. Though \\cyg\\ is also a wind-accreting system, its typical luminosity in the high state is $\\sim 0.03 L_{Edd}$ \\citep{cygx1-gierlinski-1999}, which is much lower than \\xone's. This implies that the black hole in LMC~X--1 accretes at a higher rate than its counterpart in \\cyg, perhaps because the companions' wind speeds  and densities differ, or that accretion mechanisms with different radiative efficiencies are taking place. In addition, \\cyg\\ exhibits well-documented transitions between the high/soft state and the low/hard state \\citep{cygx1-agrawal-1972, cygx1-baity-1973, cygx1-heise-1975, cygx1-holt-1976, cygx1-holt-1975, cygx1-sanford-1975}. LMC~X--1, in contrast, displays persistent soft-state emission \\citep{lmcx1-and-lmcx3-nowak-2000,lmcx1-and-lmcx3-wilms-2001}. While both \\xone\\ and \\cyg's softest states can be fitted with a power law plus disk blackbody model, they also differ markedly.  The traditional soft state model, in which the disk blackbody dominates the energy spectrum, accurately fits \\xone's spectra \\citep{ebisawa-1989, yao-2005}. \\cyg's spectra, however, are energetically dominated by the power law in both the hard and soft states. Thus, while these two systems have similar companions and orbital properties, they differ in their spectral characteristics and evolution.   An even more interesting comparison can be made between \\xone\\ and \\xthree, a black-hole binary that accretes via Roche-lobe overflow at comparable luminosity, as a means of highlighting the observational differences between disk and wind accretion. Much of this paper will be devoted to pointing out and interpreting these differences.  \\xthree\\  is a persistently bright x-ray binary in the Large Magellanic Cloud, and is the only other system with a dynamically confirmed black hole.  As such, it provides a useful counterpoint to the wind-accreting black hole binary systems mentioned above. LMC~X--3 is a $9 M_{\\odot}$ black hole \\citep{lmcx3-vanderKlis-1985} with an evolved B5 companion \\citep{lmcx3-soria-2001}.  A detailed listing of the system's dynamical properties, along with those of \\xone\\ and \\cyg, is provided in Table \\ref{tb:all-systems-orbital-params}. Differences between the systems are expected to reflect differences in the accretion flows resulting from wind versus Roche-lobe overflow accretion: LMC~X--3 accretes high angular momentum material through its first Lagrange point, while LMC~X--1 accretes low angular momentum material from its companion's stellar wind. LMC~X--3 and LMC~X--1 share spectral signatures of a persistent accretion disk, though LMC~X--1's disk spectrum defies the $L \\propto T^4$ relation \\citep{diskbb2, dunn-2010}.  \\begin{table} \\centering \\begin{minipage}{0.99\\textwidth} \\caption{System parameters for LMC~X--1, LMC~X--3, and Cygnus~X--1, with references in parenthesis. The Eddington luminosity for each source was calculated from black hole mass.} \\begin{tabular}{|p{20mm}|r|r|r|} \\hline & \\textbf{\\xone} & \\textbf{\\xthree} & \\textbf{\\cyg}\\\\ \\cline{1-4} \\hline Distance (kpc)&$48.10 \\pm 2.22$ (1) & $52.0 \\pm 0.6$ (2) & $2.5$ (3)\\\\ \\hline Black Hole Mass ($M_{\\sun}$)&$10.91 \\pm 1.41 $ (1)& $11.1 \\pm 1.4$ (4)  &$10.1$ (5)\\\\ \\hline System Inclination (degrees)& $36.38 \\pm 1.92^{\\circ}$ (1)& $50.0^{\\circ} \\leq i \\leq 70.0^{\\circ}$ (12) & $36^{\\circ} < i < 67^{\\circ}$ (11)\\\\ \\hline Companion Type& O7 III (8) & B5 IV (6)& O9.7 Iab (5) \\\\ \\hline Companion Mass ($M_{\\sun}$)& $31.79 \\pm 3.48$ (1) & $4.0 \\leq M \\leq 4.7$ (6) & $17.8$ (5)\\\\ \\hline Companion Radius ($R_{\\sun}$)& $17.0 \\pm 0.8$ (1)  & $4.4$ (6)& $22.7 \\pm 2.3$ (7) \\\\ \\hline Orbital Period (days)&$3.90917 \\pm 0.00005$ (1) &$1.70479 \\pm 0.00004$ (5)  & $5.6$ (3) \\\\ \\hline Average Soft State Luminosity (\\% L$_{Edd}$) &$\\sim 16\\%$ (8) &$\\leq 30 \\%$ (9)  & $\\sim 3\\%$ (10)\\\\ \\hline \\end{tabular}  \\label{tb:all-systems-orbital-params}  (1) \\citet{lmcx1-orosz-2009}; (2) \\citet{dibenedetto-1997}; (3) \\citet{cygx1-paczynski-1974}; (4) \\citet{gierlinski-2001}; (5) \\citet{herrero-1995}; (6) \\citet{lmcx3-soria-2001}; (7) \\citet{cygx1-companion-radius}; (8) \\citet{lmcx1-gou-2009}; (9) \\citet{lmcx1-and-lmcx3-nowak-2000}; (10) \\citet{cygx1-gierlinski-1999}; (11) \\citet{davis-1983}; (12) \\citet{lmcx3-vanderKlis-1985}.  \\end{minipage} \\end{table}   This paper examines how \\xone's spectral properties and evolution fit with, and possibly extend, existing black hole binary accretion models. In particular, our results ultimately indicate that \\xone's spectral behavior is consistent with sporadic obscuration of the innermost part of a stable accretion disk.   Section 2 of this paper presents the data processing and spectral fitting techniques employed in obtaining the results presented in Section 3.  Section 4 discusses how \\xone's unique spectral properties can be reconciled with existing accretion flow models, while Section 5 presents our conclusions and possible avenues for further investigation. ",
        "conclusions": "LMC~X--1's apparent defiance of the Stefan-Boltzmann relation \\response{may stem} from variations in the disk itself, or from variations in the scattering medium surrounding it.  Because the total number of photons is more nearly constant than the disk luminosity (see Figure \\ref{fig:lmcx1-4-panel-plot}), variations in the number of photons diverted to the power law by the corona seem more likely to explain the anomalous temperature/luminosity relations than variations within the disk itself.   The first step in quantitatively distinguishing between these two explanations for the anomalous disk relation is to compare the observed disk variation timescales to timescales inherent to the disk and to the corona.  In Figure \\ref{fig:lmcx1-long-chopped-obs-with-our-data-as-background}, \\xone\\ 's disk randomly samples a wide swath of the temperature parameter space in the space of a few hours. Individual points in the archival observation are separated by an average of 2320~s, which provides an estimate of the timescale for apparent disk variations.  We can compare this value to the disk's viscous timescale to assess whether the disk's structure is capable of changing over such a brief interval of time.   The disk's viscous timescale scales with both radius and viscosity, as \\citep{frank-king-and-raine} \\begin{equation} t_{visc} \\sim \\frac{R^2}{\\nu} \\label{eq:viscous-timescale} \\end{equation} The circularization radius, $R_{circ}$, of the accreted wind material gives a lower limit for the outer disk radius.  The disk radius, and therefore the viscous timescale, may actually be larger due to outward angular momentum transfer.  Since we are primarily concerned with finding a lower limit on the viscous timescale, however, $R_{circ}$ provides a conservative value of the disk radius. The circularization radius is defined as the radius at which the accreted matter's angular momentum due to Keplerian motion equals the angular momentum it carried upon initial capture by the black hole. Following \\citet{frank-king-and-raine} we assume that the wind's initial angular momentum about the black hole goes as \\begin{equation} l \\sim \\frac{1}{4}\\omega R_{acc}^2 \\label{eq:orbital_ang_mom} \\end{equation} where $R_{acc}$ represents the black hole's gravitational capture radius, and $\\omega$ is the black hole's orbital angular velocity about the companion. From there, we obtain \\begin{equation} R_{circ} = \\frac{G^3 M_{BH}^3 \\omega^2}{v_{rel}^8} \\textrm{.} \\label{eq:Rcirc} \\end{equation} where $M_{BH}$ represents mass of the black hole, and \\begin{equation} v_{rel}^2 = v_w^2 + v_{orb}^2 \\textrm{.} \\label{eq:vrel} \\end{equation} Therefore, the disk's size is strongly dependent on the speed of the stellar wind, which undergoes line-driven acceleration \\citep{CAK, winds-owocki-1994, winds-kudritzki-2000}. The wind's velocity increases with distance from the star \\citep{lamers-1999, kudritzki-1989} as \\begin{equation} v_w(r) = v_{\\infty} \\left[1 - 0.9983 \\frac{R_*}{r} \\right]^{\\beta} \\textrm{.} \\label{eq:vwind} \\end{equation} Here, $ v_{\\infty}$ is the wind's terminal velocity and $R_*$ is the companion's radius.  The factor of $0.9983$ is chosen such that $v_{w}$ equals the escape velocity at the star's surface.  The values of $v_{\\infty}$ and $\\beta$ are poorly constrained for O-star winds, but typically range from 1000--2000 km/s and from 0.5--1.5, respectively \\citep{kudritzki-1989},\\citep{lamers-1999}.  We conservatively assume $v_{\\infty} = 1700$ km/s, in order to compare our results with other works that consider clumping in O-star winds \\citep{the-italians}. We carry the effects of the full range of $\\beta$ values through all subsequent calculations to obtain a range of results.   The final step in calculating $R_{circ}$ is determining $v_w$ where the wind reaches the accretion cylinder. To find $R_{acc}$, we equate the wind's gravitational and kinetic energy terms with respect to the black hole and find \\begin{equation} R_{acc} = \\frac{2GM_{BH}}{v_{rel}^2} \\textrm{.} \\label{eq:Racc} \\end{equation} Solving (\\ref{eq:vrel}), (\\ref{eq:vwind}) and (\\ref{eq:Racc}) numerically over the full range of $\\beta$ values constrains $v_w$ to lie between $154$ and $635$ km/s by the time it reaches the black hole's accretion cylinder.   $R_{circ}$ is thus limited to values between $6.7 \\times 10^8$ and $2.4 \\times 10^{10}$ cm, in agreement with previous work \\citep{lmcx1-and-lmcx3-nowak-2000}.   Having found $R_{acc}$, we turn to calculating the disk's viscosity, $\\nu$.  The Shakura-Sunyaev $\\alpha$-disk prescription posits \\begin{equation} \\nu = \\alpha c_s H \\label{eq:alpha-prescription} \\end{equation} where $c_s$ and $H$ are the disk's local sound speed and height, respectively. Observations of transient x-ray binary outbursts indicate $\\alpha \\sim 0.1$ \\citep{alpha-disk-values}. Like \\xone, x-ray transients while in outburst host ionized accretion disks around a central black hole. These similarities motivate us to use $\\alpha = 0.1$ in the following calculations. Substituting equilibrium thin-disk formulae \\citep{frank-king-and-raine} for the sound speed and disk height into (\\ref{eq:alpha-prescription}) gives \\begin{equation} \\nu = 1.8 \\times 10^{14} \\alpha^{4/5}\\dot{M}_{16}^{3/10} m_1^{-1/4} R_{10}^{3/4} \\left[1 - \\left(\\frac{R_{min}}{R} \\right)^{1/2} \\right]^{3/10} \\textrm{ cm}^2\\textrm{/g} \\textrm{.} \\label{eq:alpha-disk-viscosity-formula} \\end{equation} Here, $m_1$ is $M_{BH}$ in units of $M_{\\sun}$, $\\dot{M}_{16}$ is the accretion rate in units of $10^{16}$ grams per second, $R_{10}$ is radius in units of $10^{10}$ cm, and $R_{min}$ is radius of the innermost stable orbit. Simple geometry requires that the accretion rate depends on the companion's mass loss rate as \\begin{equation} \\dot{M}_{BH} = \\frac{1}{4} \\left(\\frac{R_{acc}}{R_{sys}}\\right)^2 \\dot{M}_{*} \\textrm{ g/s} \\label{eq:MdotBH} \\end{equation} with a typical O-star $\\dot{M}_*$ of $10^{-5}$ $M_{\\sun}$ per year. The factor of four in the denominator reflects the the fact that the black hole presents a circular cross-section to the spherically expanding wind. Equation (\\ref{eq:MdotBH}) shows that the accretion rate onto the black hole ranges between $2.2 \\times 10^{17}$ and $9.6 \\times 10^{18}$ g/s, depending on the chosen value of $\\beta$.  These values are in good agreement with accretion rates derived from the observed x-ray luminosity using \\begin{equation} L \\sim 0.1 c^2 \\dot{M}_{BH}, \\label{eq:lum} \\end{equation} which range between $8.4 \\times 10^{17}$ and $2.0 \\times 10^{18}$ g/s. The $0.1$ efficiency factor in (\\ref{eq:lum}) is recommended by \\citet{frank-king-and-raine}. Equation \\ref{eq:lum} arises from the condition that all gravitational potential energy up to the last stable orbit is radiated away.   Using (\\ref{eq:viscous-timescale}) to combine the results of (\\ref{eq:alpha-disk-viscosity-formula}) and (\\ref{eq:Rcirc}) gives $5\\times 10^4 \\le t_{visc} \\le 4\\times 10^6$ seconds.  These values of $t_{visc}$ are at least an order of magnitude longer than the $\\sim 10^3$ second timescale of inner disk variation indicated by Figure \\ref{fig:lmcx1-long-chopped-obs-with-our-data-as-background}. This provides the first piece of evidence that the observed disk luminosity variations are not caused by a process inherent to the disk itself. But the calculation of the disk viscous timescale above assumes that the entire disk is involved in any change. The Roche-lobe accreting black-hole binary GRS~1915+105 shows very fast spectral changes in many of its variability states near the Eddington luminosity \\citep{grs1915-belloni-2000}, which appear to be repeated ejections of the innermost parts of the disk.  But when the inner disk is ejected and the spectrum turns hard in GRS~1915+105, the total x-ray count rate also drops dramatically.  It is the absence of any such change in count rate in \\xone\\ that leads us to favor obscuration by marginally optically thick, ionized coronal material to an ejection mechanism.   The slope of the temperature-luminosity relation that the chopped archival observations in Figure \\ref{fig:lmcx1-long-chopped-obs-with-our-data-as-background} occupy provides the second indication that \\xone's apparent disk luminosity variations are caused by a process unrelated to the accretion disk. If the variations in disk luminosity were due to changes in the accretion rate, the points would move sequentially along the Stefan-Boltzmann relation shown in Figures \\ref{fig:lmcx3-disktemp-vs-lum-for-flash} and \\ref{fig:lmcx1-disktemp-vs-lum-for-flash}.    Since the anomalous disk temperature-luminosity relation is poorly explained by changes within the accretion disk, we ask whether changes in the corona could explain the relation.  First, we consider which coronal properties are compatible with the observations.  Then, we speculate how such a corona might form.   Figure \\ref{fig:lmcx1-scattering-frac-vs-disk-temp} provides a clue as to how the coronal material factors in \\xone's anomalous disk temperature-luminosity relation. The fraction of disk photons that suffer inverse Compton scattering while traversing the corona is lowest when the apparent inner disk temperature is around $0.9$ keV. This value lies at the higher end of the observed range of disk temperatures for this source.  The inverse Compton scattering fraction increases with decreasing inner disk temperature, up to a scattering fraction of nearly unity at $T_{in}$ of $\\sim 0.4$ keV. Such behavior is consistent with a steady disk whose inner radii are sometimes obscured by a cloud of energetic electrons.  In that case, the number of disk photons, which form the seed photons for inverse Compton scattering in the overlying corona, would remain constant. The fraction of original photons that suffer up-scattering would depend only on the amount of coronal material present. Figure \\ref{fig:lmcx1-4-panel-plot} shows that \\xone's total photon flux remains largely constant, while its coronal luminosity varies substantially.  To quantify this model, we integrate the emission from the $\\alpha-$disk flux prescription \\citep{frank-king-and-raine} outwards from inner radii at temperatures ranging between $0.4$ and $0.9$ keV.  Figure \\ref{fig:lmcx1-disk-lum-frac-change} shows the resulting temperature-luminosity relation, as fitted to all of the chopped archival datasets with four or more sub-intervals. The disk temperature and luminosity are constant in each observation.  The amount of inner disk obscuration is the only variable that determines where the chopped observations fall along the best-fit line.   The model appears consistent with all the data, particularly with observation 30087-01-04, which provides the strongest constraint.  Note that these individual chopped observations run roughly perpendicular to the Stefan-Boltzmann relation in temperature-luminosity space.  Suzaku or NuSTAR observations may allow the detection of variations of the power-law index on relevant (hour) timescales; hardness variations might then be interpreted as variations in Comptonizing optical depth, in which case hardness might be expected to correlate with the flux in the tail.   Figures \\ref{fig:lmcx1-disk-lum-frac-change} and \\ref{fig:lmcx1-scattering-frac-vs-disk-temp} raise questions about the analysis in \\citet{lmcx1-gou-2009}, which derives the spin of LMC~X--1's black hole from measurements of the inner disk temperature. Their analysis requires that the inner disk always be unobscured. Our analysis, however, shows that the degree of inner disk obscuration in LMC~X--1 varies significantly on short timescales.  In systems such as LMC~X--3, where the disk is not obscured, there is no difficulty with the approach of \\citet{lmcx1-gou-2009}.  \\begin{figure} \\includegraphics[width=0.9\\textwidth]{figure9.eps} \\caption{The obscured inner disk model, fitted to all of the chopped archival observations containing more than three sub-intervals. Observation numbers are listed in the upper left corner of each plot.  Black points represent the chopped observations, while gray points represent our data.  The solid black line plots the obscured inner disk model, which has been normalized to fit the chopped observations. \\label{fig:lmcx1-disk-lum-frac-change}} \\end{figure}  Why, in fact, does LMC~X--3 have less apparent inner disk obscuration than LMC~X--1 does?   We can imagine \\response{at least} three possible answers. The first option is that the coronal optical depth is much larger, in general, in LMC~X--1 than in LMC~X--3.  This possibility is ruled out by the overlap between the two sources in terms of the flux ratio between the disk and power-law components;  when the disk luminosity in \\xthree is in decline, a significant power-law tail can form in that system without disrupting the excellent Stefan-Boltzmann relation for the disk component \\citep{smith-2007}. A second possible answer supposes that the geometries of the coronae are very different: LMC~X-3 has a geometrically large corona, while LMC~X--1 has one that is more centrally concentrated, so that it obscures more of the inner disk while producing the same amount of upscattering. A third possibility supposes that the corona in both cases has a conelike shape (possibly a jet, but not necessarily, from our evidence alone), and that the apparent differences between the two sources are due to an inclination effect.  In \\xone, which has an inclination of $36.38\\pm 1.92^{\\circ}$ \\citep{lmcx1-orosz-2009}, the conical corona may block the inner parts of the disk as we look down through it, while in LMC~X--3, with an inclination between $50^{\\circ}$ and $70^{\\circ}$ \\citep{lmcx3-vanderKlis-1985}, the cone presents itself in semi-profile, where it doesn't obscure the inner disk but we can still see the Compton upscattered flux it produces coming out its sides.  While either of the latter two pictures is consistent with our data, the last is more appealing \\response{for two reasons. Firstly}, it does not require an unexplained difference in the form of the corona in the two binaries. \\response{Secondly, it may someday be possible to verify or to disprove using x-ray polarization observations. A 'down the barrel' jet observation (e.g. LMC~X--1) would sport a lower over-all scattering polarization fraction than an observation of the same type of geometry viewed 'in profile' (e.g. LMC~X--3) would. The precise magnitude of the expected polarization fraction difference will govern how easily its presence or absence can be verified. In any case, future observations with the GEMS x-ray polarimeter will at least begin to explore the relevant parameter space, provided the instrument can spend upwards of $10^5$ seconds observing each source \\citep{GEMS-paper}.}    Finally, we turn to the question of how a population of energetic electrons forms and dissipates above the innermost region of the disk on timescales shorter than $t_{visc}$. Figure \\ref{fig:wind-accretion-geometry-cartoon} illustrates one possible accretion scenario that could account for LMC~X--1's unusual properties.  In this picture, the incoming wind material sheds its net angular momentum in the ``post shock'' before accreting onto the black hole \\citep{hoyle-lyttleton-accretion}. Most of the shocked gas remains in the disk plane, and enters the accretion disk (middle arrow). A smaller portion of the heated post-shock matter launches on randomly distributed radial trajectories towards the black hole (upper and lower arrows).  This component of the accretion flow rapidly heats on its radial descent into the black hole's gravitational potential well.  Unlike the disk, which is dense enough to cool efficiently via bremsstrahlung, the diffuse infalling matter cannot efficiently cool via free-free interactions.  By the time the radially accreting matter passes over the inner disk, it can have reached electron temperatures above the $100$ keV required for the observed inverse Compton upscattering \\citep{narayan-and-yi-1995,ichimaru-1977,rees-1982,esin-1998}. This model allows the corona to change independently of the disk.   The presence of two accretion flow components, which are largely independent beyond the post-shock, is consistent with the constant observed total photon number and minimal disk luminosity variation in \\xone. Changes in the wind mass flux, due perhaps to wind clumping, register immediately in the radial, coronal flow.  Viscous diffusion in the disk, meanwhile, smooths out those same sharp mass discontinuities. While the coronal surface density, and therefore the total energy of coronal emission changes rapidly, the disk emission evolves more slowly and less dramatically. The scenario of independent flows with the thin disk remaining intact beneath the corona, even when the latter is optically thick, is consistent with the picture derived from observations of the hard state and hard-to-soft transitions in 1E 1740.7-2942 and GRS 1758-258 \\citep{smith-2002}. Recent observations of the iron fluorescence line in black holes in the hard state have shown relativistic line profiles indicating that the inner disk remains present even in the hard state \\citep{miller06}, until extremely low lumniosities, below 1\\% of Eddington, are reached \\citep{tomsick09}.  In the luminosity range of LMC~X--1, we therefore expect the inner disk to remain present, whether the thin disk is actually continuous or whether its inner portion is recondensed from a disk that is disrupted at intermediate radii \\citep{liu11}.   \\begin{figure} \\includegraphics[width=0.7\\textwidth]{figure10.eps} \\caption{A cartoon of the accretion geometry of wind accretion systems.  Incoming material from the companion's wind (dashed gray lines) must dissipate its angular momentum in a shock front behind the compact object, a.k.a. the 'post shock',  before it can be accreted (thick arrow). \\label{fig:wind-accretion-geometry-cartoon}} \\end{figure}    Differing patterns of x-ray variability (e.g. Figure~\\ref{fig:lmcx1-and-lmcx3-disk-temp-vs-date-short}) show promise as a means of identifying black hole companions when optical observations cannot. Fast changes in the inner disk temperature (Figure~\\ref{fig:lmcx1-and-lmcx3-disk-temp-vs-date-short}), and an anomalous temperature/luminosity relation (Figure~\\ref{fig:lmcx1-disktemp-vs-lum-for-flash}), both combined with a stable net photon flux (Figure~\\ref{fig:lmcx1-4-panel-plot}), may signal the presence of a wind accretor, perhaps combined with a low inclination, just as long-term hysteresis in state changes appears to be a signature of Roche-lobe overflow accretion and a large disk \\citep{smith-2002, smith-2007}.   The type of two-component accretion flow proposed in this paper is best suited to black hole binary systems where the companion drives a high-velocity stellar wind. At the same time, this particular model leaves ample room for other coronal generation mechanisms that have already been proposed for black hole binary systems accreting via Roche-lobe overflow.  At least two types of  observations can either help to confirm or to disprove the hypothesis we have proposed. The first focuses on recent x-ray observations of IC~10~X--1 and NGC~300~X--1 \\citep{ic10x1-barnard-2008}. Both systems have disk to power-law luminosity ratios, as well as high-mass main-sequence companions, that resemble \\xone's. Binary systems that contain both an O-type star capable of driving such a wind and a black hole companion are rare, which has made this particular combination of accretion mechanisms particularly hard to observe. Further observations of IC~10~X--1 and NGC~300~X--1, however, can potentially determine how deep their apparent similarities to \\xone\\ run, and will provide more opportunities to test our model's veracity. The second test relies on observations of intermediate state, low-disk-fraction black hole binaries. These systems exhibit unusual temperature-luminosity relations akin to the ones seen in \\xone\\ \\citep{dunn-2011}. One can confirm whether the same mechanism is at work in both \\xone\\ and an intermediate state system by measuring the variability (or lack thereof) of the latter's total photon flux.  Stable photon fluxes would support the proposition that these systems also contain stable, but variably-occulted, inner disks.   \\xone\\ has provided specific evidence that stellar mass black hole binaries may have x-ray spectral features which are unique to wind accretion. In a more general sense, the fact that \\xone\\ shares certain x-ray spectral properties with other types of black hole binary systems motivates further investigation of its seemingly unorthodox accretion mechanism's full scope and utility."
    },
    "1107/1107.2666_arXiv.txt": {
        "abstract": "{We suggest the existence of a fundamental connection between baryonic and dark matter. This is motivated by both the stability of these two types of matter as well as the observed similarity of their present-day densities. A unified genesis of baryonic and dark matter is natural in models in which the baryon number is promoted to a spontaneously broken local gauge symmetry. This is illustrated in a specific class of SUSY models using the Affleck-Dine mechanism. The dark matter candidate in these scenarios is charged under the baryon gauge symmetry and must have a mass around the GeV scale to give the correct present-day abundance.  We discuss constraints from B-factories, LEP, mono-jet searches at the Tevatron, and dark matter direct detection experiments.  A baryonic dark force is shown to be consistent with all data for mediators as light as the GeV scale.  } ",
        "introduction": "The stability of the proton in the standard model (SM) is the consequence of an \\emph{accidental} symmetry of the renormalizable Lagrangian. Generic UV completions of the SM are hence expected to induce proton decay. Even if ad hoc symmetries are introduced in the renormalizable formulation of these theories, the question remains: what prevents the proton from decaying through higher dimension operators? We interpret the stability of the proton as one of the few experimental facts that physics beyond the SM should explain.  Global symmetries are commonly viewed as accidental at low energies, as they are expected to be violated by quantum gravity effects. In order to guarantee the conservation of a certain global number one should therefore appeal to more fundamental, dynamical arguments.   A simple way to ensure the stability of the proton is to promote the baryon number to a local symmetry.  In this paper we explore the consequences of gauged baryon number, spontaneously broken at the weak scale. Within the framework proposed here, proton stability is ensured at all orders in perturbation theory, and is not spoiled by non-renormalizable effects.  The idea of gauging the baryon number has been studied in a number of papers in the eighties and nineties~\\cite{rajpoot,he,foot,MC1,MC2,Bailey,Carone}, and more recently in~\\cite{Wise,Wise2,Dong:2010fw,Ko:2010at,baryonbump,Perez:2011dg,Perez:2011pt,Lebed:2011fw}. The main difference between our work and those already present in the literature is that here the lepton and baryon sectors are treated in a totally \\emph{asymmetric} way, as we gauge baryon number to protect the proton stability,  but allow explicit lepton number violation.  This approach is phenomenological in spirit, since there are no observations suggesting a fundamental relationship between leptons and baryons. For example, the stability of the electron is ensured in any theory where it is the lightest electrically charged fermion. This follows from an interplay between electric charge conservation and kinematics without any need for additional dynamical assumptions. Furthermore, the smallness of the neutrino masses can be elegantly arranged via the see-saw mechanism, where lepton number is explicitly violated.   In contrast, there are indications of a connection between the baryonic sector and dark matter.  Such a connection is suggested by two observational facts: (1) both dark matter and protons are stable on cosmological timescales without any obvious explanation a priori; and (2) precise cosmological observations reveal that their present-day abundances are remarkably similar, being only a factor of about 5 apart.   Interestingly, the promotion of baryon number to a local symmetry requires the existence of a new anomalous global symmetry. We find that under rather general dynamical conditions this new physics provides a natural dark matter candidate. Moreover, the genesis of visible and dark matter in fact \\emph{unifies} in such a way that their relic abundance can be similar for comparable masses.  Our scenario overlaps with the ``asymmetric dark matter\" scenario \\cite{Kaplan,Agashe:2004bm,ADM,hooper,Hylogenesis,Darko,Dodelson,Xo,Falko,adm9,adm12,adm5,AWIMP,Frandsen:2011kt,MarchRussell:2011fi,Davoudiasl:2011fj,Cui:2011qe}, although here the asymmetry is not generated in one sector and then subsequently transferred to the other.  Rather, in our framework both asymmetries are generated simultaneously~\\cite{pangenesis,cogenesis}.  We focus on SUSY realizations, although this is not a necessary ingredient in our treatment. The main advantages offered by SUSY realizations are that (i) the theory is technically natural and calculable up to very high scales;  and (ii) the genesis of matter can be elegantly explained by the Affleck-Dine mechanism \\cite{AD,DRT,DRT2}.    The main features of our framework are as follows: \\begin{itemize}  \\item We promote baryon number to a gauge symmetry $U(1)_{B}$. Anomaly cancellation implies the existence of new exotic quarks with their own anomalous $U(1)$ global symmetry.  \\item To facilitate the decay of exotic quarks, a SM singlet chiral superfield $X$ is introduced.  The dark matter $X$ has gauge baryon interactions and is stable, since it is the lightest particle charged under the new anomalous $U(1)$.  \\item The DM plus the visible sector have a \\emph{single}, linearly realized, nonanomalous accidental global symmetry $U(1)_D$.  Baryogenesis requires a primordial asymmetry in $U(1)_D$, which necessarily implies the existence of an asymmetry in the dark sector.  \\item This general setup has supersymmetric flat directions, allowing for the {\\it simultaneous} generation of dark and baryonic asymmetries via the Affleck-Dine mechanism.  \\item The proper DM abundance is achieved via the annihilation mode to quarks mediated by the new baryonic gauge boson $Z_{B}$.  \\item The $Z_{B}$ gauge boson also mediates interactions between dark matter and nuclei.  Though the coupling to nuclei is vectorial, the coupling to DM is model dependent.    \\end{itemize}    The outline of the paper is the following.  In Section~\\ref{models} we present a class of UV complete models in which the baryon number is gauged, and discuss the phenomenological viability of these models.  Models in which the dark matter has either a purely vectorial or purely axial coupling to the $U(1)_B$ gauge boson are presented in Section \\ref{UV}.    In Section~\\ref{baryogenesis} we see that both the visible and dark sectors have comparable primordial asymmetries, and show that typical perturbative realizations of our scenarios require the DM candidate to have a mass between sub-GeV to tens of GeV.  In Section~\\ref{Bounds} we present a general analysis of the bounds arising from B-factories, LEP and Tevatron, and direct detection experiments. Here we work under the assumption that the only remnant of the dark sector are the baryon gauge field and a DM particle  carrying charged under the gauged baryon number.  If vectorially coupled to the $U(1)_B$ gauge boson, the DM must be around the GeV scale to evade direct detection constraints, whereas axial couplings lead to novel momentum and velocity suppression such that direct detection bounds are much weaker.  Constraints from mono-jet searches at the Tevatron are strong when the mediator is produced on-shell.  Our main numerical results are summarized in Fig. \\ref{Tev100-axial} for two benchmark examples. The first panel shows the allowed region for a 1 GeV dark matter particle with a purely vector coupling to the $U(1)_B$ gauge boson.  The second panel shows the allowed region for a 10 GeV dark matter particle having a purely axial coupling to the $U(1)_B$ gauge boson.  While our cosmological predictions for the ratio of the dark matter asymmetry to the baryon asymmetry are sensitive to the UV completion, our collider and direct detection bounds are independent of the specific realization for the cancellation of the $U(1)_B$ anomalies.    \\begin{figure}% \\begin{center} \\includegraphics[width=2.9in]{everything.pdf} \\includegraphics[width=2.9in]{axialeverything.pdf} \\caption{\\small Constraints on a dark matter candidate. {\\em Left}: $m_X=1$ GeV and a purely vector coupling with $q_V=4/3$. {\\em Right:} $m_X=10$ GeV and a purely axial coupling with $q_A=4/3$. The filled areas are excluded by $\\Upsilon \\rightarrow invisible$ (left) or $\\Upsilon \\rightarrow hadrons$ (right)  and the CDF monojet search. Tevatron projections assuming $10$ fb$^{-1}$ of integrated luminosity are also shown (dashed). The solid red line indicates the values for $(m_B, g_B)$ required to get the right DM abundance for a symmetric DM species (WIMP cross section). Asymmetric DM species must lie above the red curve. \\label{Tev100-axial}} \\end{center} \\end{figure} ",
        "conclusions": "In this paper we have gauged baryon number to provide a unified framework that ensures proton and dark matter stability and relates the present-day abundances of these two types of matter.  In order for the ordinary baryon number to be embedded into a gauge symmetry $U(1)_{B}$ one needs to postulate the existence of a chiral, $q'$-sector beyond the SM. A viable, perturbative formulation of this scenario also requires a dark matter particle $X$ such that processes ``$q'\\rightarrow$ SM + $X$\" are allowed, implying that experimental signatures of the $q'$ particles are different from a conventional 4th generation of fermions.  The DM in these models is generally charged under the gauge baryon number $U(1)_B$, and therefore its dominant interactions with the visible sector are with the SM quarks. We discussed the collider constraints and direct detection bounds on the new leptophobic force under the assumption that the coupling of the dark matter to the new gauge boson is either purely vectorial or purely axial. We found that for a vector coupling our models are consistent with current data provided both the DM and the $U(1)_B$ force carrier have masses at the GeV scale. For purely axial couplings the direct detection bounds are weakened considerably allowing for larger dark matter masses, while the collider constraints are essentially unchanged.    A characterizing feature of our program is that the genesis of baryons and DM in the early Universe is unified.  The DM is asymmetric and its primordial asymmetry is generally comparable to the baryon asymmetry, motivating further study of GeV scale DM."
    },
    "1107/1107.0455_arXiv.txt": {
        "abstract": "In this manuscript, the structure of broad emission line regions (BLRs) of well-mapping double-peaked emitter (AGN with broad double-peaked low-ionization emission lines) 3C390.3 is studied. Besides the best fitted results for double-peaked broad optical balmer lines of 3C390.3 by theoretical disk model, we try to find another way to further confirm the origination of double-peaked line from accretion disk. Based on the long-period observed spectra in optical band around 1995 collected from AGN WATCH project, the theoretical disk parameters of disk-like BLRs supposed by elliptical accretion disk model (Eracleous et al. 1995) have been well determined. Through the theoretical disk-like BLRs, characters of observed light-curves of broad double-peaked H$\\alpha$ of 3C390.3 can be well reproduced based on the reverberation mapping technique. Thus the accretion disk model is preferred as one better model for BLRs of 3C390.3. Furthermore, we can find that different disk parameters should lead to some different results about size of BLRs of 3C390.3 from the one measured through observational data, which indicates the measured disk parameters are significantly valid for 3C390.3. After that, the precession of theoretical elliptical disk-like BLRs being considered, we can find that the expected line profile in 2000 by theoretical model is consistent with the observed line profile by HST around 2000. Based on the results, we can further believe that the origination of broad double-peaked balmer emission lines of 3C390.3 are from accretion disk around central black hole. ",
        "introduction": "The properties of objects with broad double-peaked low-ionization emission lines (hereafter, double-peaked emitters) have been studied for more than two decades, since the double-peaked emitters were found in nearby radio galaxies (Stauffer et al. 1983, Oke 1987, Perez et al. 1988, Halpern 1990, Chen et al. 1989, Chen \\& Halpern 1989). Some theoretical models have been proposed to explain the properties of double-peaked broad emission lines which can be used as one probe of the broad-line regions of active galactic nuclei (Eracleous et al. 2009), such as binary black hole model (Begelman et al. 1980, Gaskell 1983, 1996, Boroson \\& Lauer 2009, Lauer \\& Boroson 2009, Zhang et al. 2007), double stream model (Zheng et al. 1990, 1991, Vellieux \\& Zheng 1991), accretion disk model (Chen et al. 1989, Chen \\& Halpern 1989, Eracleous et al. 1995, 1997, Bachev 1999, Hartnoll \\& Blackman 2000, 2002, Karas et al. 2001, Gezari et al. 2007, Flohic \\& Eracleous 2008, Lewis et al. 2010, Tran 2010, Chornock et al. 2010, Gaskell 2010) etc.. Although which theoretical model is more preferred for double-peaked emitters is still an open question, the accretion disk model which indicates double-peaked broad emission lines are from central accretion disk is so far the more widely accepted model applied to explain the properties of double-peaked broad emission lines.  Due to the un-obscured emitting region on the receding jet, the double-stream model can be ruled out for double-peaked emitter 3C390.3 (Livio \\& Xu 1997). Through the long-period observational results, binary black hole model has also been ruled out due to unreasonable large central BH masses for some double-peaked emitters, such as Arp102B, 3C390.3 (Eracleous et al. 1997). Furthermore, based on the long-period variations of double-peaked broad emission lines, accretion disk model is successfully applied to reproduce the characters of varying observed line profiles of double-peaked emitters, such as the individual double-peaked emitter NGC1097 (Storchi-Bergmann et al. 1995, 1997, 2003), and one sample of double-peaked emitters (Lewis et al. 2010, Gezari et al. 2007). However, we should note that although accretion disk model is so far the widely accepted theoretical model for double-peaked emitters, the other theoretical models can be expected to be preferred for some special double-peaked emitters, such as the x-shaped radio objects which could be candidates for objects with binary black holes in central regions (Merritt \\& Ekers 2002, Zhang et al. 2007). Thus, only based on the double-peaked appearance of broad emission lines, there is no confirmed way to affirm which theoretical model is more preferred for double-peaked emitter.  Among the sample of double-peaked emitters (Eracleous \\& Halpern 1994, 2003, Eracleous et al. 1995, Strateva et al. 2003), 3C390.3 is one well studied double-peaked emitter, also one well-studied mapping AGN (Dietrich et al. 1998, O'Brien et al. 1998, Leighly et al. 1997, Shapovalova et al. 2001, 2010, Sergeev et al. 2002, Peterson et al. 2004, Pronik \\& Sergeev 2007, Sambruna et al. 2009, Jovanovic et al. 2010). Based on the reverberation mapping technique (Blandford \\& Mckee 1982) and virialization method (Peterson et al. 2004, Collin et al. 2006, Peterson \\& Bentz 2006, Peterson 2010), size of BLRs (distance between central black hole and broad line emission gas clouds) and central virial BH masses of 3C390.3 have been well determined (Peterson et al. 2004, Onken et al. 2004, Kaspi et al. 2005, Bentz et al. 2006, 2009, Brandon \\& Bechtold 2007). Besides the common BH masses and size of BLRs, through the theoretical accretion disk model applied for observed double-peaked broad emission lines, some basic disk parameters of supposed theoretical disk-like BLRs of 3C390.3 have also been measured (Eracleous \\& Halpern 1994), such as the inclination angle of accretion disk, the inner and outer radius of BLRs in accretion disk etc.. However, only due to the best fitted results for double-peaked balmer lines by accretion disk model, we can not firmly confirm that the double-peaked broad emission lines are exactly coming from disk-like BLRs in accretion disk around central black hole for 3C390.3.  Besides the best fitted results for double-peaked broad emission lines by theoretical models, based on the pioneer work by Blandford \\& Mckee (1982), geometrical structures of BLRs can be mathematically structured by so-called transfer function in the reverberation mapping technique through some special mathematical methods (such as the Maximum Entropy Method, Narayan \\& Nityananda, 1986) applied to the properties of long-period observational varying both continuum emission and broad line emission (Peterson et al. 1993, 1994, Horne et al. 1991, Goad et al. 1993, Wanders \\& Horne 1994, Pijpers \\& Wanders 1994, Krolik 1994, Winge et al. 1995, Bentz et al. 2010).  However, it is unfortunate that there should be NOT unique solution to the so-called transfer function used to determine the structures of BLRs of AGN, due to the less complete information about variations of continuum and broad emission lines (Maoz 1996). Thus, even based on the solution of transfer function, the supposed theoretical and mathematical geometrical structures of BLRs of AGN can NOT still be firmly confirmed.  In the previous remarkable work about structures of BLRs of AGN, the main objective is to determine the structures of BLRs through observational properties/characters. However, we think it will be also very interesting to check whether the observational properties of long-period variations of broad line emission and/or continuum emission could be better reproduced through the reverberation mapping technique, based on the expected structures of BLRs supposed by theoretical model, which is the main objective of our paper.  The structures of our paper are as follows. Section 2 gives some information about observed spectra of 3C390.3 around 1995 and how to find theoretical disk parameters for the supposed disk-like BLRs through the best fitted results for broad double-peaked balmer lines by theoretical elliptical accretion disk model (Eracleous et al. 1995). Section 3 shows the results through the reverberation mapping technique under the disk-like BLRs for 3C390.3. Finally section 4 gives some detailed discussions and conclusion. The cosmological parameters $H_{0}=70{\\rm km\\cdot s}^{-1}{\\rm Mpc}^{-1}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{m}=0.3$ have been adopted here. ",
        "conclusions": "As one well-known double-peaked emitter, based on the disk parameters measured above, to check the expected line profile of double-peaked broad balmer emission line of 3C390.3 in different years/periods is very interesting. Although the more recent observed line profiles of double-peaked broad H$\\alpha$ of 3C390.3 in Jan. 2007 can be found in Shapovalova et al. (2010), the observed spectrum in Sep. 2000 by HST STIS is collected, because the reduced spectrum in 2000 is public and conveniently collected from the website of Multimission Archive at STScI (MAST) (http://archive.stsci.edu/) (HST Proposal 8700, PI: Prof. Andre Martel in Space Telescope Science Institute). The detailed description about observational technique for the spectrum around 2000 by HST can be found in Popovic (2003). Here, we mainly check whether the observed line profile of broad balmer emission lines in 2000 can be reproduced through theoretical elliptical accretion disk model based on the disk parameters measured from spectra observed around 1995 for double-peaked emitter 3C390.3.  Before proceeding further, the precession period should be determined. Through the obtained disk parameters, the relativistic precession period of the elliptical disk-like BLRs around central black hole of 3C390.3 can be simply calculated (Weinberg 1972), \\begin{equation} \\begin{split} T_{pre} &= \\frac{2\\times\\pi}{\\delta_{\\phi}} \\sim \\frac{2\\times\\pi}{\\frac{6\\pi\\times G\\times M_{BH}}{c^2A(1-e^2)}} \\\\ & \\sim 70 {\\rm years}\\hspace{4mm} ({\\rm inner\\ \\ radius})\\\\ & \\sim 370 {\\rm years}\\hspace{4mm} ({\\rm outer\\ \\ radius}) \\end{split} \\end{equation} where 'A' is the semi-major axis length and $e$ is the eccentricity. Although, the precession period for outer elliptical rings is very long, the short precession period for inner elliptical rings should lead to some apparent variations of double-peaked broad emission lines from 1995 to 2000.  We simply assume that orientation angle is the unique parameter to vary from 1995 to 2000 in the elliptical accretion disk model. As what we have done in section 3, the elliptical disk-like BLRs are separated into 60 rings. It is clear that orientation angle of each ring in 2000 should be some different due to different semi-major axis length for each ring, $\\phi_{2000,i}=\\phi_{1995,i}\\pm 2\\times\\pi\\times\\frac{5}{T_{pre,i}}$, where $\\pm$ means the rotating direction of disk-like BLRs is clockwise or anticlockwise for observer. After the orientation angles of the 60 rings around 2000 are determined, line profile expected in 2000 can be re-constructed by the sum of 60 expected double-peaked components from 60 rings, as what we have done in Section 3. Here the theoretical line profile in 2000 is mainly based on the mean spectrum around 1995 in AGN WATCH project. The mean observed spectrum shown as thin solid line in Figure~\\ref{2000} is created by PCA (Principal Components Analysis, so-called Karhunen-Loeve Transform) method applied for all the 66 observed spectra of broad H$\\alpha$ around 1995. The convenient and public IDL PCA program 'pca\\_solve.pro' written by D. Schlegel in Princeton University is used, which is included in the SDSS software package of IDLSPEC2D (http://spectro.princeton.edu/). Commonly, the first principal component represents the mean spectrum. From the mean spectrum, it is clear that there are similar flux ratio ($\\sim1.46$) of blue peak to red peak of broad double-peaked H$\\alpha$ and similar flux ratio ($\\sim1$) of blue part to red part with those mean values shown in Figure~\\ref{var}.  Figure \\ref{2000} shows the probable observed line profile of broad double-peaked H$\\alpha$ of 3C390.3 around 2000, and shows the comparison between the mean observed spectrum around 1995 and observed spectrum around 2000. Based on the theoretical accretion disk model, through the results shown in top-left panel in Figure~\\ref{2000}, we can find that due to precession period, after about 5 years, the flux ratio of red peak to blue peak should be changed. In the figure, flux density at blue peaks of all the line profiles have been normalized to 1. If rotating direction of disk-like BLRs is clockwise to observer, the expected flux ratio (1.36) of red peak to blue peak of spectrum around 2000 should be smaller than the one (1.46) of observed spectrum around 1995. If rotating direction of disk-like BLRs is anticlockwise to observer, the expected flux ratio (1.68) in 2000 should be larger than the one (1.46) in 1995. Through the observed spectrum in 2000, the flux ratio of blue peak to red peak is about $1.79\\pm0.18$. Thus, if anticlockwise rotating disk-like BLRs are applied for 3C390.3, the variations in observed spectra in different years can be naturally explained, although the applied model is over-simplified. The some larger value 1.79 than theoretical expected 1.68 are perhaps due to the not so accurate precession period, the probable existence of varying of hot spots and/or warped structures, different local broadening velocity (which should lead the part around red peak to be some flat), etc.. The not so consistent red peak positions for spectrum around 2000 and mean spectrum around 1995 is probable due to the variation of emissivity power slope during the 5 years from 1995 to 2000.  Finally, a simple summary is listed as follows. Based on the observed spectra around 1995 for 3C390.3, the elliptical accretion disk model is applied to fit the observed double-peaked broad H$\\alpha$ collected from AGN WATCH project. Then based on the disk parameters, the formed geometrical structures of BLRs of 3C390.3 are applied to check the reverberation mapping results for 3C390.3. The disk-like BLRs supposed by theoretical elliptical accretion disk model can well re-produce the observational results about reverberation mapping technique. Thus the disk-like BLRs are preferred for 3C390.3. Furthermore, we check the effects of precession of elliptical disk-like BLRs on the observed line profile. After about 5 years, the line profile of broad double-peaked H$\\alpha$ observed in 2000 by HST can be well expected by the theoretical elliptical accretion disk model, based on the disk parameters measured through spectra observed around 1995, which confirms the elliptical accretion disk model is appropriate to double-peaked emitter 3C390.3\t."
    },
    "1107/1107.2982_arXiv.txt": {
        "abstract": "The Hubble Constant measured from the anisotropy in the cosmic microwave background (CMB) is shown to be independent of small changes from the standard model of the redshift dependence of dark energy. Modifications of the Friedmann equation to include phantom power (w $<$ --1), textures (w = --2/3) and curvature are considered, and constraints on these dark energy contributors from supernova observations are derived. Modified values for the density of matter inferred from cosmic density perturbations and from the CMB under these circumstances are also estimated, as exemplified by 2df and SDSS. ",
        "introduction": "Turner (1999) coined the term dark energy to name the power source of the accelerating Universe (Garnavich \\etal 1998; Perlmutter \\etal 1999).  For the first time, there was a plausible, complete accounting of matter and energy in the Universe. Expressed as a fraction of the critical density there are: neutrinos, between 0.3\\% and 15\\%; stars, 0.5\\%; baryons (total), 5\\%; matter (total), 40\\%; smooth, dark energy, 60\\%; adding up to the critical density. This accounting is consistent with the inflationary prediction of a flat Universe and defines three dark-matter problems: Where are the dark baryons? What is the nonbaryonic dark matter? What is the nature of the dark energy? The leading candidate for the (optically) dark baryons is diffuse hot gas; the leading candidates for the nonbaryonic dark matter were slowly moving elementary particles left over from the earliest moments (cold dark matter), such as axions or neutralinos; the leading candidates for the dark energy involve fundamental physics and include a cosmological constant (vacuum energy), a rolling scalar field (quintessence), and light, frustrated topological defects.  Gooding \\etal (1992) considered a Universe in which density fluctuations are produced in an initally smooth Universe by the ordering dynamics of scalar fields following a symmetry breaking phase transition at the grand unified scale. Such transitions lead to the formation of an unstable topological defect known as ``global texture.\"  Carroll \\footnote{ (http://ned.ipac.caltech.edu/level5/Carroll2/Carroll2$\\_$1.html)} points out that for some purposes it is useful to pretend that the $-ka^{-2} R_0^{-2}$ term in the Friedmann equation represents an effective ``energy density in curvature'', and define $\\rho_k$ % $-(3k / 8\\pi GR_0^2) a^{-2}$.  Caldwell (1999 astro-ph 8168) remarks that most observations are consistent with models right up to the w = --1 or cosmological constant limit, and so it is natural to ask what lies on the other side at w $<$ --1. He termed this phantom energy.  In this paper we outline how a dark energy program will constrain these elements and, in particular, how they affect the measurement of the Hubble Constant by means of the anisotropy in the cosmic microwave background. Section 2 extends the Friedmann equation; $\\S$3 shows that supernova data are currently tolerant of small values of $\\Omega_1$ and $\\Omega_2$; $\\S$4 explores the degeneracies in CMB data; $\\S$5 examines how matter density experiments like 2dF (Peacock \\etal 2001) are affected; $\\S$6 broadly explores the parameter space of $\\Omega_n$ as it applies to SN and CMB data. Our conclusions are in the final section. ",
        "conclusions": "Our primary conclusion is that introducing $\\Omega_1$ or $\\Omega_{-1}$ does not change WMAP values of H$_0$ or $\\Omega_m$ . It is easy to show that this conclusion extends to phantom energy generally for w $<$ --1 with $\\Sigma a^{-n}\\Omega_n~+~a^{-x}\\Omega_x~ = ~h^2$ and x $<$ 0. Using the supernova data alone, it is not possible to determine all the $\\Omega$'s because of degeneracies. \u00a0But in combination with CMB data, the degeneracies are broken. Second, we find that $\\Omega_1$ $<$ 0.2 and $\\Omega_{-1}$ $<$ 0.1 with 95\\% confidence. Stricter limits will follow from dark energy program. Third, $\\Omega_2 ~\\approx~ -0.2\\Omega_1$ for $\\Omega_1~ <<$ 1. If $\\Omega_2~ = ~\\epsilon$ (say, 10$^{-6}$) due to inflation, $\\Omega_1~<~0.018/(f_1-f_0)~=~0.06$."
    },
    "1107/1107.5046_arXiv.txt": {
        "abstract": "The properties of dark matter halos, including mass growth, correlate with larger scale environment at fixed mass, an effect known as assembly bias.  However, whether this environmental dependence manifests itself in galaxy properties remains unclear. We apply a group-finding algorithm to Data Release 7 of the Sloan Digital Sky Survey to estimate the dark matter halo mass of each galaxy and to decompose galaxies into those that exist at the centers of distinct halos and those that orbit as satellites within larger halos. Using the 4000-\\AA\\ break as a measure of star formation history, we examine the correlation between the quenched fraction of galaxies, $\\fq$, and large-scale environment, $\\rho$. At all galaxy magnitudes, there is a positive, monotonic relationship between $\\fq$ and $\\rho$. We use the group catalog to decompose this correlation into the contribution from central and satellite galaxies as a function of halo mass. Because satellites are more likely to be quenched than central galaxies, the observed $\\fq$-$\\rho$ correlation is primarily due to variations of the halo mass function with environment, which causes a larger fraction of satellite galaxies at high $\\rho$. For high-mass central galaxies, there is a weak but positive correlation between $\\fq$ and $\\rho$ at both fixed galaxy luminosity and fixed halo mass. For low-mass central galaxies ($\\mgal\\la 10^{10.0} \\hhmsol$), there is no correlation between $\\fq$ and $\\rho$. The latter results are inconsistent with the strong assembly bias of dark matter halos seen in this mass regime if recent galaxy growth at all correlates with recent halo growth, as we demonstrate through a high-resolution $N$-body simulation. We also find that the mean stellar age of quenched central galaxies is independent of $\\rho$ at fixed $\\mgal$, while the formation times of low mass halos vary significantly. We conclude that the processes that halt the star formation of low mass central galaxies are not correlated to the formation histories of their host halos, and old galaxies do not reside preferentially in old halos. ",
        "introduction": "Many galaxy properties are correlated with their large- and small-scale environment. The primary examples are the color-density relation and the morphology-density relation (e.g., \\citealt{oemler:74, davis_geller:76, dressler:80}). These results demonstrated that, in the local Universe, galaxies in lower-density regions are typically less massive, and that even at fixed luminosity or mass, they are bluer and more likely to be spirals.  With the advent of large-scale galaxy redshift surveys such as the Sloan Digital Sky Survey (SDSS; \\citealt{york_etal:00}), these seminal works have been refined with increased statistical precision (\\citealt{hogg_etal:04, kauffmann_etal:04, blanton_etal:05a, baldry_etal:06, park_etal:07, bamford_etal:09}) and extended via direct measurements of the correlations between star formation rates and environment (\\citealt{balogh_etal:04, kauffmann_etal:04}) and via measurements out to $z\\approx 1$ (\\citealt{cucciati_etal:06, cooper_etal:07}).  Studies that investigated the relative importance of environment on different scales indicated that small-scale environment (Mpc scale) correlates stronger with galaxy properties than large-scale environment (\\citealt{kauffmann_etal:04, blanton_etal:06a, blanton_berlind:07, wilman_etal:10}). These results suggest that the mass of the dark matter halo in which a galaxy resides plays the dominant role in the galaxy's formation and evolution.  For example, \\cite{blanton_berlind:07} found that the blue fraction of galaxies in rich groups and clusters is independent of environment at fixed halo mass.  In this paper we extend these results by focusing on galaxies in low-mass halos.  The definition of environment is thus key to any discussion of this topic. In this paper as in many other works, `environment' is defined as the local \\textit{galaxy} density at some smoothing scale. However, the physical interpretation of this definition of environment is not entirely straightforward because galaxy density is a biased measure of dark matter density, in addition to being subject to shot noise and the effects of redshift-space distortions (see, e.g., \\citealt{cooper_etal:05}).  The choice of scale is also important. In this paper we make the distinction between (large-scale) environment and halo mass. We define environment as the local galaxy density on scales larger than the host dark matter halo of a given galaxy, using a fiducial smoothing scale of 10 \\hmpc. Halo mass is similar to `small-scale environment', but the distinction is important. The dispersion of host halo masses for galaxies at a fixed 1 Mpc local density can be 1 dex or more and using distance to the $n$-th nearest neighbor yields an even larger spread in halo mass \\citep{haas_etal:11}.  This scatter demonstrates the need for quantifying halo mass explicitly.  The idea that galaxy properties depend primarily on host halo properties is physically motivated, since the halo virial radius corresponds to a physical transition between infall and virialized motions, capable of supporting strong shock fronts and hot, thermalized gas \\citep[e.g.,][]{dekel_birnboim:06}.  In this sense, a galaxy group necessarily corresponds to a single, virialized dark matter halo.  This is distinct from many other definitions of `group environment'.  For example, the `Local Group' is not a group at all in this context, but rather a close pair of separate halos. The difference is an important one: for a halo to enclose both the Milky Way and Andromeda (with one galaxy at the halo center), it would have to be larger than $10^{13}$ \\msol to equal their separation of $\\sim 800$ kpc, a mass at which diffuse halo gas is heated to high temperatures, causing shock heating and stripping of the disks in our galaxy or M31 (or both). As a close pair of $\\sim 10^{12}$ \\msol\\ halos, their tidal interactions and gas physics is markedly different: there is no ram pressure and tidal stripping is much weaker.  It is thus more informative to quantify halo mass rather than galaxy density on a small scale.  However, environment does play a role in the formation of dark matter halos.  In the numerical Universe, dark matter halos exhibit `assembly bias': at fixed mass, halo properties depend on their formation histories (\\citealt{gao_etal:05, harker_etal:06, gao_white:06, wechsler_etal:06, wetzel_etal:07, li_etal:08, dalal_etal:08}). At low halo mass, older and more concentrated halos form in high density environments. At high mass the effect reverses, with younger, less concentrated halos forming in high-density regions, though the effect is also weaker.  Semi-analytic models of galaxy formation predict that this assembly bias propagates into the evolution of galaxies at fixed halo mass, especially in the formation histories and star formation properties of low-luminosity galaxies in low-mass halos (\\citealt{zhu_etal:06, croton_etal:07}). At first glance, these predictions appear in concert with the above observational evidence that the fraction of galaxies (at fixed luminosity) that are red and quenched is highest at high densities.  Another approach to understanding the connection between galaxies and dark matter is the Halo Occupation Distribution (HOD; see, e.g., \\citealt{seljak:00, peacock_smith:00, roman_etal:01, berlind_weinberg:02, cooray_sheth:02}). The HOD has emerged as our most powerful tool for understanding the clustering of galaxies and quantifying the bias between galaxies and dark matter, where bias here is defined as the ratio between the galaxy and matter densities on an arbitrary smoothing scale. The central quantity in the HOD is $P(N|M)$, the probability that a halo of mass $M$ contains $N$ galaxies within a defined sample. The galaxy sample may be arbitrarily defined, thus $P(N|M)$ will be different for every sample definition. The bias of galaxies depends strongly on galaxy properties, with strong clustering for galaxies that are brighter, redder, and more elliptical (e.g., \\citealt{norberg_etal:01, norberg_etal:02, zehavi_etal:02, zehavi_etal:05, zehavi_etal:10, li_etal:06}). In the HOD context, these trends are all explained by the different halos that different types of galaxies occupy. Brighter central galaxies live in more massive halos that are more highly clustered. At fixed galaxy mass, redder or more elliptical galaxies are more likely to be satellites in groups or clusters, thus enhancing their clustering.  These observations fit into the HOD formalism without any explicit correlation between galaxy formation and environment, other than the fact that massive halos form at high densities (e.g.~\\citealt{bond_etal:91}).  The $P(N|M)$ ansatz has been extended even further in the `subhalo abundance matching' approach, which (in its simplest form) assumes a monotonic relationship between halo mass, $M$, and galaxy luminosity, $L$ (\\citealt{kravtsov_etal:04, wang_etal:06, vale_ostriker:06, conroy_etal:06, moster_etal:09, conroy_wechsler:09, wetzel_white:10, behroozi_etal:10}).\\footnote{One can use halo maximum circular velocity rather than halo mass, or use galaxy stellar mass rather than galaxy luminosity.}  Here, if the number density of halos above $M$ is the same as that of galaxies above $L$, then $M$ halos host galaxies of $L$ brightness. They key to this approach is to include not just isolated halos, but also `subhalos', defined as bound, virialized halos located within the virial radii of larger host halos. The primary assumption of this approach is that (sub)halo mass entirely determines the luminosity of the galaxy within it, regardless of any other properties of the (sub)halo or the galaxy.   The quantity $P(N|M)$ makes the explicit, and strong, assumption that $N$ depends only on $M$.  That is, the properties of galaxies, including luminosity, color, age, and morphology, depend only on the mass of their host dark matter halo and have no additional dependence on a halo's formation history or environment.  This formulation was not proposed with any theoretical prejudice about galaxy formation, but rather it was the simplest implementation and was successful in fitting the clustering properties of nearly all available data. However, the discovery of halo assembly bias, combined with the complex relation between galaxy properties and environment, especially those governing galaxy growth and star formation, calls into question this ansatz.  In \\cite{tinker_etal:08_voids}, we sought to test this $P(N|M)$ ansatz through different galaxy clustering measures that probe different environments. We used the two-point correlation function, $\\xi_g(r)$, to constrain halo occupation and make predictions for the void probability function. The predictions were in excellent agreement with the measurements, demonstrating the validity of the HOD from high to low densities. \\cite{abbas_sheth:06} and \\cite{skibba_etal:06} used other galaxy clustering measures to demonstrate that the assumption of $P(N|M)$ is consistent with observed data for luminosity-based galaxy samples. However, \\cite{yang_etal:06} and \\cite{wang_etal:08_assembly_bias} found that the clustering of low-mass galaxy groups at fixed mass depends on the color of the galaxies contained within them such that groups with redder galaxies are more highly clustered.  In this paper, we examine the relation between galaxy star formation histories, their host dark matter halos, and their large-scale environment.  We explore in particular the extent to which the formation histories of galaxies relate to those of their host halos and whether the assembly bias seen in simulation extends to the galaxy population.  We focus on low-mass galaxies, the regime in which halo assembly bias effects are the strongest, and we examine star formation history based on narrow 4000-\\AA\\ break, $\\dn$, which, unlike color, is insensitive to dust reddening.  In order to observationally determine halo masses for SDSS galaxies and separate this from large-scale environment, we construct galaxy group catalogs using the \\cite{yang_etal:05} group-finding algorithm.  Our use of a group finder is critical because galaxies of the same properties may occupy halos of a wide range of mass. At fixed stellar mass, galaxies may live at the center of low-mass halos or they may exist as satellites in group or cluster-sized halos. Anywhere between 10-40\\% of galaxies are satellites, depending on galaxy mass and color (see, e.g., \\citealt{yang_etal:08, zehavi_etal:10}). The distinction between central and satellite galaxies will be the cornerstone of understanding the environmental correlations of galaxy properties.  The evolution of satellite galaxies is more complex, and their properties also depend on position within a halo (e.g., \\citealt{vdb_etal:08b,hansen_etal:09,weinmann_etal:10}). Here we focus on the properties of central galaxies in halos below the group scale. We will explore the observed dependence of satellite star formation rate on galaxy mass, halo mass, and halo-centric radius in detail in a companion paper (Wetzel, Tinker, \\& Conroy~2011a), hereafter referred to as Paper II.  Finally, in a third paper (Wetzel, Tinker, \\& Conroy~2011b), hereafter referred to as Paper III, we will use these measurements combined with a high-resolution simulation to test various physical mechanisms for the quenching of satellite galaxies.  We examine galaxies selected on either $r$-band magnitude or stellar mass, $\\mgal$, and we examine star formation history based on $\\dn$. Magnitude and $\\dn$ have the advantage of being well-defined observationally, though in Paper II we will focus entirely on stellar mass and instantaneous specific star formation rate (SSFR), quantities which are derived and thus more model-dependent, but also more readily physically interpretable.  We have examined all halo and environmental trends in both papers selecting on both $M_r$ and $\\mgal$ as well as examining both $\\dn$ and SSFR.  While these lead to slight quantitative differences in our results, they do not change any of our results qualitatively.  Throughout, we define a galaxy group as any set of galaxies that occupy a common dark matter halo (including cluster-mass halos), and we define a halo as having a mean interior density 200 times the background matter density. A {\\it host halo} is a halo that is distinct: its center does not reside within the radius of a larger halo. A {\\it subhalo} is one whose center is located within the radius of a larger halo. For all calculations we assume a flat, \\lcdm\\ cosmology of $(\\om,\\s8,\\omb,n_s,h_0)=(0.27, 0.82, 0.045,0.95,0.7)$. For galaxy magnitudes we do not assume a value of $h$, but for brevity we write all magnitudes as $M_r$ rather than $M_r-5\\log h_0$. Stellar masses are in units of $\\hhmsol$.   \\begin{figure*} \\centering \\begin{tabular}{cc} \\psfig{file=magr_mstellar.ps,width=0.5\\linewidth,clip=} & \\psfig{file=lumfunc.ps,width=0.5\\linewidth,clip=} \\end{tabular} \\caption{ \\label{samples} {\\it Left panel}: Relation between $r$-band magnitude, $M_r$, and stellar mass, $\\mgal$, with the latter taken from the {\\tt kcorrect} code of \\citet{blanton_roweis:07}. Vertical lines show the magnitude limits for our volume-limited samples as listed in Table 1, while horizontal lines indicate the stellar mass limits within each volume-limited sample. {\\it Right panel}: Cumulative number density of galaxies as a function of both $M_r$ and $\\mgal$, offering a rough conversion between the two. The star indicates $M_r^\\ast$ from the \\citet{blanton_etal:03} luminosity function.} \\end{figure*}    \\begin{table} \\centering \\begin{minipage}{140mm} \\caption{Volume-Limited Samples} \\begin{tabular}{@{}ccccc@{}} \\hline $M_r$ & $z_{\\rm max}$ & $N_{\\rm gal}(<M_r)$ &  $\\log M_\\ast/[h^{-2}\\,M_\\odot]$ & $N_{\\rm gal}(>M_\\ast)$\\\\ \\hline -18.0 & 0.040 & 31729 & 9.4 & 21423 \\\\ -19.0 & 0.064 & 74987 & 9.8 & 54119 \\\\ -20.0 & 0.103 & 131658 & 10.2 & 100852 \\\\ \\end{tabular} \\end{minipage} \\end{table}   \\begin{figure} \\centerline{\\psfig{figure=color_dn4k.ps,width=9.0cm}} \\caption{ \\label{color_dn4k} The distribution of $g-r$ color and $\\dn$ for $M_r=[-19,-20]$ galaxies. The shaded region shows the distribution of color, with scale on the bottom $x$-axis. The red dashed curve shows the distribution of $\\dn$ for the same galaxies, with scale on the top $x$-axis. Even though 53\\% of galaxies are red $(g-r>0.7)$, only 35\\% of galaxies are truly quenched $(\\dn>1.6)$. The blue and black dotted curves show the color distribution for quenched and active galaxies, respectively, according to their $\\dn$ values. Quenched galaxies nearly always have red colors, while 23\\% of active galaxies have red colors, primarily because of dust extinction (e.g., \\citealt{maller_etal:09}).} \\end{figure}   \\begin{figure} \\centerline{\\psfig{figure=galprop_histo.ps,width=8.0cm}} \\caption{ \\label{galprop_histo} Distribution of $\\dn$ values for galaxies in magnitude bins. Faint galaxies are predominantly active while bright galaxies are predominantly quenched but still have a significant active population. All sample distributions are bimodal with a minimum at $\\dn \\approx 1.6$. We refer to galaxies with $\\dn > 1.6$ as `quenched'.} \\end{figure}   \\begin{figure} \\centerline{\\psfig{figure=galprop_density.ps,width=8.0cm}} \\caption{ \\label{galprop_density} Quenched fraction, $\\fq$, defined by $\\dn > 1.6$, as a function of large-scale (10 \\hmpc) density of galaxies. At $\\rhob>1$, $\\fq$ rises monotonically with density, while in underdense regions the quenched fraction is nearly constant.} \\end{figure} ",
        "conclusions": "\\subsection{The Environmental Dependence of Galaxy Properties}  In this paper we have demonstrated that the properties of galaxies at fixed halo mass are independent of density measured on scales larger than the halo radius, and that the dependence of the quenched fraction of galaxies on large-scale environment can be explained by the variations of the halo mass function with environment and by disentangling the relative contributions of central and satellite galaxies.  Our results extends earlier works that demonstrated that galaxy properties correlate weakly with large-scale environment when small-scale environment is fixed on a 1 Mpc scale (\\citealt{kauffmann_etal:04,blanton_etal:06a}).  However, we have argued that, because different galaxies live in different halos, it is not possible (or physically motivated) to find a single smoothing scale that fully determines galaxy properties at all masses.  For rich groups and clusters, our results are in good agreement with \\cite{blanton_berlind:07}, who found that the the blue fraction of galaxies is independent of density at fixed group luminosity, though we do see evidence for some change with density for the brightest central galaxies.  These results are in good agreement with the results of \\cite{peng_etal:11}, who also find that the environmental dependence of the red fraction of galaxies is driven primarily by satellite galaxies and not by centrals. In detail, the observations diverge from this scenario for quenched bright central galaxies and in the mean stellar ages of central galaxies on the star-forming sequence. But our results support the notion that halo mass is the dominant quantity governing the properties of the galaxies contained within them.  \\subsection{Halo Assembly Bias and the $\\fq$-Density Relation}  As discussed in \\S 1, the formation histories of dark matter halos depend on environment. For low-mass halos, old halos form in high-density environments. For high-mass halos, old halos form in low-density environments, although there is some debate on the amplitude of this correlation. If halo and galaxy formation histories are correlated, then our measurements are inconsistent with both of these theoretical results. Figs.~\\ref{fq_den_decomp} and \\ref{mdot_delta3} clearly demonstrate a lack of correlation between $\\fq$ and $\\rhonorm$ for low-mass galaxies. Fig.~\\ref{dn4k_age} demonstrates that the mean stellar ages of quenched central galaxies also is constant with $\\rhonorm$. For galaxies below the knee is the stellar mass function, the ansatz that old galaxies live in old halos is not supported by the data.  However, the quenched fraction of central galaxies in $\\mhost\\sim 10^{12.3}$ \\hmsol\\ halos show a positive correlation with environment. This is the mass scale at which halos in simulations show little-to-no correlation between age and $\\rhonorm$. One possibility is that the environmental dependence for high-mass central galaxies arises from those that have passed through the virial radius of a more massive halo (and therefore reside higher density environments) but are counted as centrals at the present time \\citep{gill_etal:05,ludlow_etal:09,wang_etal:09}.  However, these ejected satellites are expected to become less common at increasing galaxy mass, so the effect is opposite to the trend we see here (we will discuss these processes in more detail in Papers II and III). Another possibility is that our assumption that the central galaxy is the most massive galaxy is becoming less valid at higher galaxy mass, as suggested by \\citet{skibba_etal:11}, and we are confusing high-mass centrals and satellites.  However, in order to masquerade as the signal in Fig.~\\ref{color_density_centrals}, this effect would have to have strong environmental dependence, which remains unclear. Additionally, at higher halo masses, central and satellite galaxies have increasingly similar quenched fractions. Thus, the stellar mass dependences of \\cite{skibba_etal:11} and central/satellite quenched fractions should cancel out at some level.   Our results have mixed agreement with those of \\cite{wang_etal:08_assembly_bias}, who used a similar group-finding algorithm. They found that galaxy groups with redder overall colors have higher clustering than groups of the same mass but with bluer galaxies, in agreement with our results.  However, they also found that the effect is stronger in less massive groups, opposite to the mass dependence we see here. Wang et.~al.~use color as a proxy for age as opposed to (dust-insensitive) $\\dn$, though it is unclear why using color would lead to a stronger signal.     \\subsection{The Stellar to Halo Mass Ratios of Low-Mass Galaxies}  Galaxies that are star forming and galaxies that are quenched form in halos that growth the same average rate. This implies that the stellar mass to halo mass ratio, $\\mgal/\\mhalo$ will evolve differently for quenched and star-forming galaxies. Fig.~\\ref{sham} demonstrates a $\\mhalo=10^{11.5}$ \\hmsol\\ halo that contains a central galaxy quenched at $z=0.6$ will have a $\\mgal/\\mhalo$ ratio lower by a factor of 2.5 at $z=0$ relative to the mean mean for that halo mass. However, weak lensing and satellite kinematical measurements of the halos masses of low-mass galaxies find that the halo masses are the same within the errors, regardless of color selection (\\citealt{mandelbaum_etal:06_gals, more_etal:10}). This is in conflict with the conclusion that galaxy growth and halo growth are uncorrelated. There are several possible resolutions to this discrepancy.  \\begin{enumerate} \\item{The observations of halo masses of low-mass central galaxies are biased due the use of color rather than $\\dn$ for sample segregation. The low-mass red samples will contain many intrinsically star-forming galaxies, thus bringing the mean halo masses closer together. Lack of significant numbers also increase observational errors and low masses.} \\item{Low-mass galaxies on the red sequence were quenched recently, and thus have not had enough time to create a significant difference in the $\\mgal/\\mhalo$ ratio. There is a clear red sequence at $z=1$, but the fraction of quenched central galaxies at $\\mgal\\la 10^{10}$ is not known.} \\item{Sub-$\\lstar$ central galaxies do not follow a one-way path to the red sequence; galaxies migrate back and forth from the blue cloud to the red sequence before permanently becoming quenched. Some blue galaxies at $z=0$ may have been temporarily quenched previously, while some $z=1$ red galaxies may still grow through star formation later.  This would minimize the difference between the halo masses of red and blue central galaxies at fixed $\\mgal$.} \\end{enumerate}  Galaxies that are quenched, for example, by AGN heating or major galaxy mergers (e.g., \\citealt{croton_etal:06a, bower_etal:06, hopkins_etal:08b, somerville_etal:08}) could still accrete gas or merge with gas-rich objects at some later time, replenishing their fuel supply and moving them back into the star-forming sequence. In hydrodynamical simulations of galaxy formation, low-mass halos ($\\mhalo\\lesssim 10^{12}$ \\hmsol) accrete their gas in a cold state (\\citealt{keres_etal:05, dekel_birnboim:06, keres_etal:09, brooks_etal:09}). Thus previous merging or AGN activity that might temporarily quench a galaxy will not prevent any new gas from migrating to the center of the potential well within a dynamical time, even if the stellar population is old. However, none of these solutions are mutually exclusive, and it is possible that a all three cause the measurements presented here.   \\subsection{What Quenches Low-mass Central Galaxies?}  The fraction of quenched central galaxies in the deepest voids is the same as that in regions around clusters.  Thus the mechanism by which low-mass central galaxies stop their star-formation must be environmentally independent. \\cite{wang_etal:09_dwarfs} note that many red dwarf galaxies are isolated, but propose that the majority of quenched central galaxies of this type are ejected satellites, low-mass halos that have passed through a group or cluster. The uniformity of $\\fqcen$ for low-mass galaxies argues strongly against such a model due to the strong environmental dependence that would result from such a model. We cannot rule out two mechanisms that operate primarily at low and high density, respectively, that add to a constant $\\fqcen$, but Occam's razor argues against such a scenario.  We find that the Sersic indices of these galaxies does not vary with density, with $\\langle n_s\\rangle\\approx 3.5$. This implies that the mechanism that quenches low-mass centrals also induces morphological transformation. It also strengthens the claim that the process that halts star formation in centrals at high densities also operates at low densities.  Using semi-analytic models, \\cite{croton_farrar:08} find a small population of central galaxies quenched by AGN heating, but this process can only effect high-mass halos of $\\mhalo>10^{12.5}$ \\hmsol, roughly an order or magnitude more massive than the host halos of the galaxies probed here, and even more massive than the dwarf galaxies in \\cite{wang_etal:09_dwarfs}.  Halo mergers vary only weakly with large-scale dark matter density (\\citealt{fakhouri_ma:09}), implying that the galaxy merger rate would also show no environmental dependence. Major mergers are also likely to cause a transition from disk-dominated to spheroidal-dominated morphology (see, e.g., \\citealt{hopkins_etal:08b}). But it remains unclear whether there are enough mergers to fully account for these objects. We leave this to a future study (Wetzel et.~al., in preparation).  Further analysis of the morphologies and other properties of these galaxies, as well as their redshift evolution, will shed light on the mechanisms by which quenching occurs in these systems, and whether or not quenching is transient.   \\begin{figure} \\centerline{\\psfig{figure=mdot_slopes.ps,width=8.0cm}} \\caption{ \\label{mdot_slopes} The ratio of the central quenched fraction between high- and low-density environments as a function of stellar mass. Filled circles show results from the group catalog, while curves indicate results using the growth rates of corresponding dark matter halos as measured over different redshift baselines. } \\end{figure}  \\begin{figure*} \\centerline{\\psfig{figure=dn4k_age.ps,width=18.0cm}} \\caption{ \\label{dn4k_age} Panel (a): $\\dn$ values for quiescent central galaxies as a function of $\\rhonorm$. The connected dots indicate the mean $\\dn$ for different $\\mgal$ bins. The shaded regions indicate the dispersion for each stellar mass bin. Panel (b): The change in the mean quenching timescale inferred from $\\langle \\dn\\rangle$ in panel (a). Timescales are calculated from stellar population synthesis models described in the text. Here we plot $\\Delta t$ rather than absolute age because it is less sensitive to model assumptions. Connected dots represent the same stellar mass bins as in (a). The errorbars indicate the dispersion in $\\Delta t$ implied by the variance in $\\dn$. Black curves represent results from dark matter halos: the solid black curve is the change in the mean formation time $t_{85}$ for halos of mass $10^{11.5}$ \\hmsol (equivalent to the halo mass for $\\mgal=10^{9.7} \\hhmsol$ central galaxies). Dashed black curve shows $\\Delta t$ for the oldest 20\\% of halos at each density bin (equivalent to $\\fq$ for central galaxies with $\\mgal=10^{9.7} \\hhmsol$). The dotted curve shows $\\Delta t$ for halos based on $t_{1/2}$. Panel (c): same as (a), but now for star-forming central galaxies (defined as $\\dn<1.6$). Panel (d): Same as (b) but again for star-forming galaxies. In this panel, the dashed curve indicates $\\Delta t$ for the 80\\% youngest halos at each density (corresponding to the fraction of star-forming central galaxies). Note the change in $y$-axis range.} \\end{figure*}  \\begin{figure} \\centerline{\\psfig{figure=sham.ps,width=8.0cm}} \\caption{ \\label{sham} The difference in logarithmic stellar mass between the $z=0$ mean $\\mgal$-$\\mhalo$ relation and the relation for galaxies that were quenched at redshift $z$, as a function of $z$. This plot assumes that halo growth and galaxy growth are uncorrelated, thus when a galaxy is quenched the halo it resides in continues to grow at the mean rate up to $z=0$. Solid and dashed curves are for halos with $z=0$ masses of $3\\times 10^{11}$ and $10^{12}$ \\hmsol, which host central galaxies of $\\log\\mgal=9.7$ and 10.2 \\hmsol, respectively. } \\end{figure}     \\vspace{1.0cm}  \\noindent The authors would like to thank Frank van den Bosch for supplying the Yang et.~al. group catalog, as well as useful discussions on the results presented in this paper.   \\appendix"
    },
    "1107/1107.4984_arXiv.txt": {
        "abstract": "We investigate the quantitative constraint on the triple-$\\alpha$ reaction rate based on stellar evolution theory, motivated by the recent significant revision of the rate proposed by nuclear physics calculations. Targeted stellar models were computed in order to investigate the impact of that rate in the mass range of $0.8 \\leq M / \\msun \\leq 25$ and in the metallicity range between $Z = 0$ and $Z = 0.02$. The revised rate has a significant impact on the evolution of low- and intermediate-mass stars, while its influence on the evolution of massive stars ($ M \\ga 10 \\msun$) is minimal. We find that employing the revised rate suppresses helium shell flashes on AGB phase for stars in the initial mass range $0.8 \\leq M / \\msun \\leq 6$, which is contradictory to what is observed. The absence of helium shell flashes is due to the weak temperature dependence of the revised \\triplea\\ cross section at the temperature involved. In our models, it is suggested that the temperature dependence of the cross section should have at least $\\nu > 10$ at $T = \\pow{1 \\hyp 1.2}{8}$ K where the cross section is proportional to $T^{\\nu}$. We also derive the helium ignition curve to estimate the maximum cross section to retain the low-mass first red giants. The semi-analytically derived ignition curves suggest that the reaction rate should be less than $\\sim 10^{-29}$ cm$^{6}$ s$^{-1}$ mole$^{-2}$ at $\\approx 10^{7.8}$ K, which corresponds to about three orders of magnitude larger than that of the NACRE compilation. In an effort to compromise with the revised rates, we calculate and analyze models with enhanced CNO cycle reaction rates to increase the maximum luminosity of the first giant branch. However, it is impossible to reach the typical RGB tip luminosity even if all the reaction rates related to CNO cycles are enhanced by more than ten orders of magnitude. ",
        "introduction": "Nuclear reaction rates are fundamental to the modeling of stars. Among them, helium-burning reactions by the triple-$\\alpha$ process play a critical role in evolution of red giants. The cross section of the triple-$\\alpha$ reaction is largely influenced by resonances in the energy states of $\\nucm{4}{He}$, $\\nucm{8}{Be}$ and $\\nucm{12}{C}$. These resonances were proposed by \\citet{Opik1951} and \\citet{Salpeter1952}; the resonance level of \\nucm{12}{C} was predicted by \\citet{Hoyle1953} and \\citet{Hoyle1954} to be at $7.68$ MeV, confirmed experimentally with a value of $7.68 \\pm 0.03$ MeV \\citep{Dunbar1953}. The contribution by non-resonant reaction was also taken into account by \\citet{Nomoto1985} to investigate the impact on helium ignition in accreting white dwarfs and neutron stars. Their result significantly increases the triple-$\\alpha$ reaction rate by more than 20 orders of magnitude at $T \\lesssim \\pow{3}{7}$ K. However, this drastic change of nuclear reaction rates do not affect on the main characteristics of stellar evolution because the rate is not so significantly modified within the temperature range for helium ignition in stars ($T \\gtrsim \\pow{6}{7}$ K). The only one exception is the behavior of the evolution of low-mass metal-free stars \\citep{Fujimoto1990,Suda2007}. Without CNO elements initially, the hydrogen flash is triggered by CNO cycles as a consequence of the accumulation of carbon in the hydrogen-burning domain through the \\triplea. However, this phenomenon cannot currently be observed since the probability of detecting this event is significantly small from the viewpoint of both timescale and statistics. Thus the reaction rate by \\citet{Nomoto1985} has been adopted by the later compilations of nuclear reaction rates \\citep{Caughlan1988,Angulo1999} in place of previous work, which did not include the contribution of non-resonant reactions \\citep{Fowler1975}.  Recently, \\citet{Ogata2009} (OKK, hereafter) presented results on the cross section of the triple-$\\alpha$ reaction using a direct calculation of Schr\\\"odinger equation. They also provide a tremendously larger cross section at $10^{7 \\hyp 8}$ K compared with the previous compilations for this nuclear reaction rate. The difference amounts to 26 orders of magnitude around $10^7$ K compared with the NACRE compilation \\citep{Angulo1999}, i.e., the reaction rates by \\citet{Nomoto1985}, and also with the latest reaction rate \\citep{Fynbo2005}. This result gives rise to a severe inconsistency with the current understandings of the observations. In particular, low-mass red giants are no longer produced using the revised triple-$\\alpha$ reaction rate due to the premature ignition of helium at the stellar center upon arrival at the base of the red giant branch \\citep{Dotter2009}. In addition to this report, \\citet{Saruwatari2010}, \\citet{Matsuo2011}, and \\citet{Peng2010} investigate the impact on accreting C-O white dwarfs and neutron stars. For the accretion onto white dwarfs, it is suggested that the energy released during off-center helium detonation attained when employing the OKK rate dominates over that of the central carbon ignition that is responsible for Type Ia supernovae. On the other hand, it is reported that accretion onto neutron stars with the OKK \\triplea\\ rate results in burst energies from ultracompact X-ray binaries that are too low to account for those observed. \\citet{Matsuo2011} insist that the OKK rate should be reduced by a factor of $10^{2-3}$ for $10^{8} {\\rm K}$ to be consistent with the observations of X-ray bursts. Application of the OKK rate to Cepheid models by \\citet{Morel2010} also fails to explain the accumulations of observations reporting the period change, depletion of lithium and reduction of the C/N ratio, although they insist that the OKK rate may reduce the mass discrepancy between stellar evolution models and pulsation theory.  Stimulated by this revision of the \\triplea\\ rate, new calculations of \\triplea\\ rates such as a method solving the Faddeev three-body equations have been ongoing (S.~Ishikawa, private communication). It is also proposed that broad resonances other than the Hoyle state exist, which may also affect the \\triplea\\ at low temperatures \\citep[see, e.g.,][]{Kato2010}. On the other hand, \\citet{Ogata2009} point out that the treatment of non-resonant reactions at low-temperatures adopted by \\citet{Nomoto1985} is not appropriate. Therefore, it is desired to explore the reaction rates both theoretically and experimentally, and hence, it is likely that the cross section for \\triplea\\ will be modified by future works. As a calibration test for the modification of the reaction rate, it is useful to rely on stellar evolution models by analyzing the dependence of the theoretical models on the reaction rate, especially below $\\pow{3}{8}$ K where the influence is important for the evolution of low- and intermediate-mass stars.  In this paper, we provide quantitative constraints on the \\triplea\\ cross section based on stellar evolution theory. Firstly, we investigate the impact of changing the triple-$\\alpha$ reaction rate on the evolution of low- and intermediate-mass stars during the horizontal branch and AGB phase, extending our study to the evolutionary characteristics of massive stars. The possible effects of changing the \\triplea\\ rate are analyzed by comparing our models with the observed extremely metal-poor stars in the Galaxy using the SAGA database \\citep{Suda2008}. Secondly, we derive the possible range of the \\triplea\\ rate by estimating the ignition point of helium burning on the first red giant branch (FGB) in a semi-analytic method. We also explore the possibility of varying the input physics with the OKK \\triplea\\ rate adopted. In particular, we discuss the change of the reaction rates of CNO cycles.  The paper is organized as follows. In the next section, the stellar evolution models used in this work are described. The impact of changing the \\triplea\\ rates is also discussed. \\S~\\ref{sec:ignition} is devoted to the derivation of helium ignition curve and the comparison with observations. In \\S~\\ref{sec:cnocycle}, we look for a compromise with the OKK \\triplea\\ rate by varying other input physics. Summary and conclusions follow in \\S~\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  We investigated the impact of changing the reaction rate of the \\triplea\\ on stellar structure and evolution. Models with the reaction rate provided by \\citet{Ogata2009} (the OKK rate) are computed for stars in the mass range from $0.8$ to $25 \\msun$. It is shown that the OKK rate brings about several serious problems in astrophysical contexts including the disappearance of the first red giant branch (FGB) as already pointed out by \\citet{Dotter2009}. Here we have found another significant problem that low-mass stars do not experience helium shell flashes due to the weak temperature dependence of the reaction rate. The absence of helium shell flashes results in the absence of the change of surface chemical composition through the third dredge-ups. This would suppress important channels to the formation of carbon-enhanced stars since, for a majority of carbon stars, their abundances are best explained by AGB evolution or by binary mass transfer from AGB stars. In addition, it would also prevent the formation and transportation mechanisms of \\sit-process elements in low- and intermediate-mass stars. As a consequence, AGB stars would no longer be the progenitors of carbon stars and \\sit-process element enhanced stars. In particular for metal-poor halo stars, the large fraction of observed carbon-enhanced stars with \\sit-process element enhancement could not be explained by mass transfer from the former AGB stars. Observationally, the detection of Technetium in the surface of AGB stars is the supporting evidence of the dredge-up of \\sit-process elements by the deepening of the convective envelope. On the other hand, for massive stars with $M \\gtrsim 10 \\msun$, the effect on the evolution is small because the helium ignition temperature is so high that the difference in the reaction rate between the NACRE \\citep{Angulo1999} and the OKK rate is not significant.  We also quantitatively estimated the allowed range of the triple-$\\alpha$ cross section under the currently adopted input physics in stellar evolution code. In our analysis of the helium ignition point of the star, it is suggested that the \\triplea\\ rate should not be larger than $\\sim 10^{-29}$ cm$^{6}$ s$^{-1}$ mole$^{-2}$ at $10^{7.8}$ K in order for stars to still have a red giant phase. In addition, it is shown that the occurrence of helium shell flashes requires the temperature dependence of the \\triplea\\ cross section with at least $\\nu \\gtrsim 10$ at $T \\sim \\pow{1 \\hyp 1.2}{8}$ K where the cross section is proportional to $T^{\\nu}$. For other temperature ranges, the shape of the reaction rate can be arbitrary as a function of temperature only if it intersects the ignition cross section given in eq.(\\ref{eq:cs}) at $\\log T \\gtrsim 7.9$ where helium ignition occurs at the FGB luminosity of $L \\gtrsim 10^{3} \\lsun$.  A way to compensate for the changes due to the OKK rate was sought by pursuing the increase of the CNO cycle reaction rates, but found to be impossible. We looked for a solution to increase the FGB tip luminosity by increasing the CNO cycle reaction rates to accelerate the growth of helium core. Even if we adopt the enhancement factor of $> 10^{10}$ for all the reaction rates related to CNO cycles, the maximum luminosity of the FGB cannot reach $10^{3} \\lsun$.  The method to derive the \\triplea\\ rate proposed by \\citet{Nomoto1985} is severely criticized by \\citet{Ogata2009}, while the result obtained by \\citet{Ogata2009} is still far from being consistent with observations. Since the \\triplea\\ is in the wide range of applications such as accretion events on compact stars like novae, X-ray bursts and the determination of cosmological distances through the Cepheid variables, more investigations for this important reaction are desired."
    },
    "1107/1107.2931_arXiv.txt": {
        "abstract": "The `Eagle' galaxy at a redshift of 0.77 is studied with the Oxford Short Wavelength Integral Field Spectrograph (SWIFT) and multi-wavelength data from the All-wavelength Extended Groth strip International Survey (AEGIS).  It was chosen from AEGIS because of the bright and extended emission in its slit spectrum.  Three dimensional kinematic maps of the Eagle reveal a gradient in velocity dispersion which spans $35-75\\pm10$ km s$^{-1}$ and a rotation velocity of $25 \\pm 5$ km s$^{-1}$ uncorrected for inclination.  {\\it Hubble Space Telescope} images suggest it is close to face-on.  In comparison with galaxies from AEGIS at similar redshifts, the Eagle is extremely bright and blue in the rest-frame optical, highly star-forming, dominated by unobscured star-formation, and has a low metallicity for its size. This is consistent with its selection.  The Eagle is likely undergoing a major merger and is caught in the early stage of a star-burst when it has not yet experienced metal enrichment or formed the mass of dust typically found in star-forming galaxies. ",
        "introduction": "Large redshift surveys at $0 <z<1.2$ have revealed that the population of galaxies is divided into red and blue rest-frame colours, and have measured the evolution of these populations \\citep[e.g.,][]{stra,bell,fab7}. In particular, blue galaxies in the past were brighter by $B \\simeq 1.3$ mag \\citep[e.g.,][]{bell, cnaw}, more highly star-forming \\citep[e.g.,][]{kai}, and more morphologically irregular \\citep[e.g.,][]{abra, oesc}. In addition, blue galaxies have an increasing contribution to their kinematics from disordered motions the further one looks back in time \\citep{wei1, kas7}.  These disordered motions, quantified by an integrated velocity dispersion, appear to play an important role in galaxy kinematics \\citep[e.g.,][]{wei2, kas7, cres, covi, fors9, puec10, lemo1, lemo2}. A study of $\\sim 550$ galaxies with slit spectroscopy over $0.1<z<1.2$ has shown that, for galaxies with stellar masses greater than $10^{10}$ M$_{\\odot}$, there is increasing scatter in the Tully-Fisher relation (which relates the rotation velocities of galaxies to their stellar masses or magnitudes) to $z=1.2$ \\citep{kas7}.    This scatter is dominated by galaxies with disturbed morphologies \\citep{kas7}.  A large scatter in the Tully-Fisher relation is also found for galaxies in integral field spectrograph (IFS) studies at $z\\sim 0.6$ \\citep{flor,puec8,puec10a} and $z \\sim 2-3$ \\citep[e.g.,][]{law, cres, lemo2}.  When a kinematic estimator which incorporates both rotation velocity ($V$) and integrated velocity dispersion ($\\sigma$) is adopted, {\\small $S_{K} \\equiv \\sqrt{K V^2 + \\sigma^2}$} with K=0.5 \\citep{wei1}, the resulting relation with stellar mass has remarkably small scatter at all redshifts \\citep{kas7, cres, covi, puec10, lemo1, lemo2}.  This is such that disturbed galaxies have higher values of $\\sigma$ than normal disc-like galaxies.  Undisturbed disc galaxies in the local Universe have $\\sigma$ values that range from 10--35 km s$^{-1}$ and an average value of $\\sim$20--25 km s$^{-1}$ \\citep{ghasp} which are due to the relative motions of individual gas clouds in spiral arms or a thick disc.  The higher $\\sigma$ values found for many high redshift galaxies likely represent effective velocity dispersions caused by the blurring of velocity gradients on scales at or below the seeing limit \\citep{wei1, kas7}.  These may not even have a preferred plane.  The nature of these kinematically peculiar objects cannot be fully determined from slit spectroscopy alone since it is unable to probe the full 3D kinematics of galaxies. 3D integral field observations can give further clues as to whether the high $\\sigma$ values are driven by phenomena such as rotation, major or minor merger activity, and/or violent star-formation processes.  The advent of IFSs on large telescopes has allowed for detailed studies of 3D galaxy kinematics over $0.4 \\la z \\la 3$ \\citep[e.g.,][]{fors6, fors9, puec7, yang, vans, wrig, law, epin, lemo1, lemo2}. These studies have all found galaxies with large $\\sigma$ values and attribute it to star-formation and major and minor merger activity, when active galactic nuclei are not present.  However, these samples are still small ($\\la 100$) and are typically too bright to be representative of typical galaxies, although there are exceptions \\citep[e.g.,][]{puec8, yang, law}.  A representative sample of 68 galaxies at redshifts $0.4 < z < 0.74$ has been studied with the GIRAFFE IFS with the Intermediate MAss Galaxy Evolution Sequence (IMAGES) Survey. Two main results of this survey are  the discovery of a low fraction of rotating disc galaxies \\citep{flor, neic,yang} and the finding that more distant galaxies have smaller ratios of $V$ to $\\sigma$ than galaxies at lower redshifts \\citep{puec7}.  These results are consistent with the findings of slit based studies of larger samples of $\\sim$500--1000 galaxies \\citep{wei1,wei2,kas7}. Studies  of IMAGES galaxies generally attribute disturbed kinematics to major mergers.  It is highly desirable to study more galaxies with IFSs at $z\\sim1$, an epoch probed by large redshift surveys but with few IFS observations. In this paper, IFS observations of the 3D kinematics of a galaxy at $z=0.7686$ with a high $\\sigma$ value are presented along with a suite of multi-wavelength data from the All-wavelength Extended Groth strip International Survey \\citep[AEGIS,][]{aegis}. These data allow for the galaxy to be understood in terms of the general population of galaxies at its epoch and offer insight into a likely cause for its high $\\sigma$ value. A $\\Lambda$CDM cosmology is adopted throughout: $H=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$. All logarithms are base 10, and all magnitudes are on the AB system. ",
        "conclusions": ""
    },
    "1107/1107.5385_arXiv.txt": {
        "abstract": "The {\\it AKARI} All-Sky Catalogues are an important infrared astronomical database for next-generation astronomy that take over the {\\it IRAS} catalog. We have developed an online service, AKARI Catalogue Archive Server (AKARI-CAS), for astronomers. The service includes useful and attractive search tools and  visual tools.  One of the new features of AKARI-CAS is cached SIMBAD/NED entries, which can match AKARI catalogs with other catalogs stored in SIMBAD or NED. To allow advanced queries to the databases, direct input of SQL is also supported. In those queries, fast dynamic cross-identification between registered catalogs is a remarkable feature. In addition, multiwavelength quick-look images are displayed in the visualization tools, which will increase the value of the service.  In the construction of our service, we considered a wide variety of astronomers' requirements. As a result of our discussion, we concluded that supporting users' SQL submissions is the best solution for the requirements. Therefore, we implemented an RDBMS layer so that it covered important facilities including the whole processing of tables. We found that PostgreSQL is the best open-source RDBMS products for such purpose, and we wrote codes for both simple and advanced searches into the SQL stored functions. To implement such stored functions for fast radial search and cross-identification with minimum cost, we applied a simple technique that is not based on dividing celestial sphere such as HTM or HEALPix. In contrast, the Web application layer became compact, and was written in simple procedural PHP codes. In total, our system realizes cost-effective maintenance and enhancements. ",
        "introduction": "With the progress of technology, recent astronomical surveys produce huge object catalogs: e.g., Sloan Digital Sky Survey \\citep[SDSS;][]{yor00}, Two Micron All Sky Survey \\citep[2MASS;][]{skr06}, UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{law07}, etc. It is not realistic for many astronomers to handle and analyze a number of catalogs in their local computers. Therefore, services to access the data via network have to be developed and could be used for a wide variety of astronomers' requirements. In addition, it is desired that the services have application program interfaces (APIs) to develop services, since it is not efficient that all service providers have all observational data.  In modern Web-based database solutions for object catalogs --- for example, NASA/IPAC Infrared Science Archive % \\citep[IRSA;][]{ber00}, VizieR \\citep{och00}, and Virtual Observatory \\citep[VO;][]{sza01} --- portal sites are famous services. One of their common features is that we can use rich graphical user interfaces (GUIs) and a high-level programming interface based on VO standards for our various demands. On the other hand, one of the most advanced services, % SkyServer of SDSS \\citep[][]{tha04}, established a new basic model of a modern Web-based database solution for a huge catalog. SkyServer provides not only basic search tools such as radial search but also advanced search tools such as the SQL search, CAS jobs, etc. Compared with the former services, SkyServer of SDSS is developed % with focus on the availability of flexible SQL-based programming interface, rather than supporting rich GUIs and high-level interfaces. Such SQL-based programming interfaces will become very important for online services dedicated to long-life survey database, since astronomers often need advanced search or their own programs to analyze catalogs and observational data for various unique science themes.  Our archive server AKARI-CAS% \\footnt{AKARI-CAS has been developed based on open source softwares. If readers are interested in our system, we can show our source codes. Please contact us about them.} has been developed to provide Web-based access tools for AKARI All-Sky Catalogues. The tools have been developed following the concept of SDSS SkyServer, and users can search, match up, and browse stored data using the provided tools. In our service, there are tools based on our own ideas and implementation techniques. Reporting them in this article will be worthwhile information that is applicable to developments of future astronomical Web-based database systems.  In this article, we mainly discuss the implementation and design of a Web-based service for astronomical catalogs. In addition, our article will be helpful to develop basic search tools with minimum cost and to set up RDBMS security when the system allows users direct input of SQL statements. This article is organized as follows: In \\S \\ref{overview_catalogs}, we briefly introduce AKARI All-Sky Catalogues. In \\S \\ref{overview_akari-cas}, we describe an overview of our catalog archive service. In \\S \\ref{overview_technical}, we discuss the basic direction of our technical design. Technical reports for RDBMS and Web application layers are given in \\S \\ref{implementation_rdbms} and \\S \\ref{implementation_webapp}, respectively. A summary is given in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We developed AKARI Catalogue Archive Server (AKARI-CAS) to provide basic tools to access AKARI-related catalogs for various studies. Our tools have been developed following the concept of SDSS SkyServer, and users can search, match up, and browse stored data using our attractive tools.  We discussed our direction of implementation for our service that supports many levels of users. We concluded that the whole processing of tables should be implemented into an RDBMS layer. We found that PostgreSQL is the best open-source RDBMS product for our requirements, and we developed various facilities using stored functions of PL/pgSQL, C, and SQL in the database. In contrast, the Web application layer became compact; this minimizes the cost of maintenance in the long term.  We presented our simple techniques % without HTM or HEALPix to perform fast radial search and dynamic cross-identification using RDBMS. With our report of the security configuration of PostgreSQL, the information in this article will be helpful to develop future astronomical Web-based database systems. In the development of the Web application, we showed our idea to create extensive services using only standard technology and without having a huge amount of data at a datacenter. Some thought put into the definition of Web APIs and public release of the APIs can increase the possibility of online astronomical services."
    },
    "1107/1107.5666_arXiv.txt": {
        "abstract": "We cross match the NVSS and FIRST surveys with three large photometric catalogues of luminous red galaxies (LRGs) to define radio-loud samples. These have median redshifts 0.35, 0.55 and 0.68 and, by matching rest-frame optical and radio properties, we construct uniform samples across the three surveys. This paper is concerned with the clustering properties of these samples derived from the angular correlation function. The primary aim is to characterise any evolution in the clustering amplitude of radio galaxies bellow $z\\sim0.68$.  We find no evidence for evolution in the large-scale ($\\sim1-50\\,h^{-1}$\\,Mpc) clustering amplitude. Our radio galaxy autocorrelations are consistent with previous findings indicating little-to-no evolution in the redshift range 0.68 to 0 ($\\sim6$\\,Gyr of time). We also cross correlate radio galaxies with the parent LRG samples to increase the precision of our results and again find no evidence for evolution in comoving coordinates. Our results are inconsistent with a long-lived model for the clustering evolution that assumes radio sources randomly sample the LRG population. A model where the halo mass is constant with redshift is consistent with the data. This is similar to QSOs that have clustering amplitudes consistent with a single halo mass at all redshifts. Given that the brightest radio sources show stronger evolution in space density compared to fainter radio sources we restrict our samples to include only objects with $L>10^{26}$\\,W/Hz and repeat the analysis. Again we find no evidence for evolution in the comoving correlation amplitude. These radio sources appear to inhabit the same mass halos as fainter radio galaxies ($\\sim9\\times10^{13}h^{-1}$\\,M$_\\odot$). These halos are $\\sim$twice as massive as those of the general LRG population ($\\sim4\\times10^{13}h^{-1}$\\,M$_\\odot$) and $\\sim30$ times as massive as optical AGN/QSOs ($\\sim3\\times10^{12}h^{-1}$\\,M$_\\odot$). ",
        "introduction": "Radio galaxies have long been used to study large scale structure as they are highly biased indicators of the mass distribution and are naturally found in the high redshift ($z\\sim1$) Universe. Understanding the observed bias of radio galaxies requires relating their clustering to that of their host populations. Furthermore, any differential evolution in their clustering with respect to their hosts gives insight into the nature of radio galaxies and the mechanisms that trigger them.     The typical host galaxy of a radio source depends strongly on its luminosity. At lower radio luminosities ($L\\lesssim10^{23}$\\,W/Hz) the population is almost exclusively made up of starforming galaxies with related diffuse radio emission, above $L\\sim10^{23}$\\,W/Hz almost all radio sources are associated with active galactic nuclei (AGN). The hosts of the brightest radio galaxies have strong emission lines (e.g. \\citealt{h+l79,mcc93}) while lower-luminosity radio AGN are hosted by massive ellipticals with little-to-no optical line emission. This transition implies a cutoff luminosity between objects with and without emission lines that \\citet{joh08} found to be around $\\sim10^{26}$\\,W/Hz. This luminosity is close to the traditional cutoff between morphologically classified FRI and FRII objects \\citep{f+r74}. Furthermore, while the bright end of the radio luminosity function evolves strongly towards higher space densities at high $z$ \\citep{lon66,d+p90}, fainter radio AGN ($L\\lesssim10^{26}$\\,W/Hz) show much less evidence for evolution \\citep{c+j04,sad07}. This evolution in space density is roughly matched by the typical hosts of radio AGN. Massive elliptical galaxies (luminous red galaxies; LRGs), that host lower luminosity radio AGN, show little evolution in number density for $z<1$ \\citep{bel04,wak06,bro07}. Higher luminosity radio sources are more associated with optical/Xray AGN \\citep{mcc93} that show strong evolution out to $z\\sim2.5$ \\citep{sch68,croom09b}.  Radio sources are strongly biased tracers of the underlying mass distribution of the Universe \\citep{ymp89,p+n91,b+w02,mag04,bra05}. Since the advent of deep all sky surveys (e.g. FIRST, NVSS, SUMSS, WENSS) there have been major leaps in the ability to describe the environments of radio galaxies. These milli-Jansky surveys sample the luminosity range $\\sim10^{23}$ to $10^{26}$\\,W/Hz at moderate redshifts, where radio sources are found almost exclusively in LRGs. LRGs are known to be a strongly biased tracer of the matter density, and their bias increases with the mass/luminosity of the LRG.  Radio sources tend to be found in the most massive elliptical galaxies and more massive galaxies are more likely to host them \\citep{bes05,m+s07}. Recent studies have compared the clustering of radio galaxies to that of radio-quiet galaxies that have been matched in their optical properties to the radio sample. Most find that radio galaxies are significantly more clustered than optically identical samples of quiescent galaxies \\citep{wak08a,man09,don09}, although see \\citet{hic09} for a counter example we will discuss later. The indication is that the environment of a galaxy contributes to the probability of it hosting a radio source.  In this paper we are primarily concerned with how the clustering of radio sources evolves for $z\\lesssim0.7$. In this redshift interval the clustering amplitude of LRGs is approximately constant \\citep{wak08b,saw09}. However, the clustering amplitude of quasars evolves such that, when models of gravitational collapse are assumed, the implied dark halo mass for quasars is approximately constant \\citep{croom05,ros09}. We will evaluate the clustering amplitude of radio LRGs in a series of samples from $z\\sim0.68$ to 0.35 to investigate any evolution in their clustering. This will be done both for a sample of medium luminosity ($L\\sim10^{24.7}$\\,W/Hz) radio galaxies, and a high luminosity subsample ($L>10^{26}$\\,W/Hz).  In section~\\ref{sec:data} we introduce the data that will be used in this work, section ~\\ref{sec:rad_match} describes the algorithm we develop for identifying radio LRGs, section~\\ref{sec:cor_anal} describes the correlation analysis to be used, section~\\ref{sec:rlc} compares the clustering of radio-loud and quiet LRGs, section~\\ref{sec:evol} then looks at the evolution of radio LRGs and in section~\\ref{sec:Levol} we investigate the evolution of the very brightest radio LRGs. 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. All radio flux densities and luminosities are at 1.4\\,GHz unless otherwise stated and when estimating radio luminosities throughout this paper $k$-corrections have been performed assuming a continuum shape of $S_\\nu\\propto\\nu^{-0.7}$. ",
        "conclusions": "We have cross-matched three large photometric samples of LRGs to the FIRST and NVSS surveys to define three samples of radio LRGs. We have then measured the 2-point correlation function to investigate the clustering properties of the samples. Similar to previous authors we find evidence that: 1) radio LRGs are more strongly clustered than non-radio galaxies matched in terms of luminosity and colour; 2) the clustering strength of radio LRGs at $r<1$h$^{-1}$Mpc increases with radio luminosity up to $L\\sim 10^{25}$ W/Hz and then remains roughly constant but at $r>1$h$^{-1}$Mpc the clustering is independent of radio luminosity. We further find that inconsistencies between the radio-LRG cross-correlation and auto-correlation amplitudes may suggest that a simple linear bias model may be insufficient to describe the relation between radio-LRG and/or LRG clustering and mass clustering.  However, the primary goal of this work was to investigate any evolution in the clustering strength of radio galaxies with redshift. We find no evidence for evolution in the large scale ($r>1$h$^{-1}$Mpc) clustering amplitude of radio-LRGs. We show that our radio $\\times$ radio LRG autocorrelations are consistent with previous authors indicating no evolution in comoving coordinates in the redshift range $0<z<0.68$ ($\\sim6$Gyr). Furthermore, we make use of cross-correlations to increase the precision of our results and still find no evidence for evolution. We then restrict our samples to objects with $L > 10^{26}$W/Hz and again find no evidence for evolution in the correlation amplitude. Our results are consistent with a single dark halo mass of $9\\times10^{13}h^{-1}$\\,M$_\\odot$ for all radio LRGs in the redshift range $0<z<0.68$. This could be compared to  quasars that appear to inhabit halos of $m_{DH}=3\\times10^{12}h^{-1}$\\,M$_\\odot$ again reasonably independent of redshift and luminosity. The significantly more massive halos for the high-luminosity radio galaxies, at least,  may provide a problem for unified AGN models. Finally, our radio-LRG cross-correlations are inconsistent with a model in which LRG clustering follows a long-lived model and radio sources are randomly sampling the LRGs."
    },
    "1107/1107.2270_arXiv.txt": {
        "abstract": "We study the evolution of non-Gaussianity in multiple-field inflationary models, focusing on three fundamental questions: (a) How is the sign and peak magnitude of the non-linearity parameter $\\fnl$ related to generic features in the inflationary potential? (b) How sensitive is $\\fnl$ to the process by which an adiabatic limit is reached, where the curvature perturbation becomes conserved? (c) For a given model, what is the appropriate tool -- analytic or numerical -- to calculate $\\fnl$ at the adiabatic limit? We summarise recent results obtained by the authors and further elucidate them by considering an inflection point model. ",
        "introduction": "An important feature of canonical single field models is that the curvature perturbation produced at horizon crossing is conserved,\\cite{Lyth,zetaAndZetaCon} with statistics indistinguishable from Gaussian.\\cite{Maldacena:2002vr} In multiple field models, on the other hand, isocurvature modes may also be produced, and can subsequently source the evolution of the curvature perturbation as the field space path curves.\\cite{GarciaBellido:1995qq,Gordon:2000hv} The curvature perturbation and its statistics (such as the power spectrum and non-Gaussianity\\cite{Vernizzi:2006ve,otherMethods}) can therefore continuously evolve after horizon crossing, leading to far richer behaviour than that allowed in single field models. Here we focus on non-Gaussianity, and in particular the nonlinearity parameter, $\\fnl$, though many of our general conclusions extend to other observables as well.  In canonical multi-field models, the non-Gaussianity present at horizon crossing is negligible.\\cite{Seery:2005gb} In this setting, however, it is possible for $\\fnl$ to evolve to large values.\\cite{hybrid,Byrnes:2008wi,Kim:2010ud,Peterson:2010mv} To calculate the observationally relevant value of non-Gaussianity, in principle one has to follow the evolution until the time of last scattering, where the Cosmic Microwave Background (CMB) was imprinted. Given our present ignorance about the detailed physics of the early universe, this would not be possible in practice. In many models, however, a regime is reached during the evolution long before this time -- the so-called {\\it adiabatic limit} -- where all the isocurvature modes decay, and the curvature perturbation becomes conserved.   There are different ways in which such a limit could be reached, and this has consequences both for the possible values of the observable parameters, such as $\\fnl$, and for the techniques which can be reliably employed to calculate observables, i.e. whether analytic methods will suffice, or numerical methods are necessary. A useful classification of models according to when the adiabatic limit is attained is: \\begin{itemize} \\item{\\it Models in which an adiabatic limit is reached `naturally' by convergence into a `focusing region' of the potential. This can be further sub-divided into cases where the convergence occurs during slow-roll, and cases in which convergence occurs only after the slow-roll approximation fails.} \\vskip 0.1in \\item {\\it Models where an adiabatic limit is reached abruptly due to an additional degree of freedom, such as a waterfall field being destabilised.} \\vskip 0.1in \\item {\\it Models for which no focusing region in the inflationary potential exists, and an adiabatic limit can be reached only by embedding the inflationary model into a larger scenario, perhaps one which includes perturbative reheating.} \\end{itemize}   Recently an extended study of these possibilities was undertaken by the authors in the context of of multi-field models of inflation, capable of producing large non-Gaussianities.\\cite{EMST-11} Here we give a summary of those results and further elucidate them by considering an inflection point model, which is of the first type, and which can produce a large positive or negative value of $\\fnl$ at the adiabatic limit.   The structure of the paper is as follows. In \\S \\ref{background} we give a summary of the background theory which is used to formulate analytic expressions for observables, discussed in \\S \\ref{sec:analytics}, together with the conditions needed for the non-Gaussianity to be large at the adiabatic limit when it is reached naturally. In \\S \\ref{sec:largeNG} we briefly discuss features in the potential which generate large transitory non-Gaussianities during the super-horizon evolution, which may be relevant for the final observable value at the adiabatic limit if this limit is reached abruptly. Finally, we discuss the usefulness of our results in \\S \\ref{sec:models} when applied to specific models, and demonstrate this by considering a new example. We conclude in \\S \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the evolution of non-Gaussianity in multiple-field inflationary models, using analytical and numerical methods. We have shown that the descent of fields from a ridge or their convergence into a valley can result in significant growth of non-Gaussianity, with the two cases being distinguished by the sign of $\\fnl$.   To concretely predict non-Gaussianities one can employ analytical expressions or numerical methods. Currently, however, the former rely on the slow-roll approximation which limits their applicability. In such cases we have demonstrated that numerical methods can be invaluable.   To calculate observationally relevant values of parameters such as $\\fnl$, in practice it is necessary that an adiabatic limit is reached where these parameters become constant. In order to determine the possible behaviours of $\\fnl$ as it reaches this limit, and the techniques needed to calculate it accurately there, we have found it is useful to classify inflationary models according to whether the adiabatic limit is reached naturally by the convergence of field space trajectories during the slow-roll regime, naturally after slow-roll ends, abruptly, or where no focusing region exists, only after reheating occurs. We have summarised the recent results from our paper\\cite{EMST-11} concerning these classes of models and given a new illustration here using a new model.   Finally, we note that the new example included in this work indicates that it is easy to construct two-field sum--separable models which exhibit large $\\fnl$ of positive or negative sign, even when the adiabatic limit is reached naturally during slow-roll inflation. We note, however, that all such models exhibiting a large non-Gaussianity at the adiabatic limit appear to require a large degree of fine tuning in their initial conditions."
    },
    "1107/1107.5295_arXiv.txt": {
        "abstract": "Although pulsars are one of the most stable clocks in the universe, many of them are observed to 'glitch', i.e. to suddenly increase their spin frequency $\\nu$ with fractional increases that range from  $\\frac{\\Delta\\nu}{\\nu}\\approx 10^{-11}$ to $10^{-5}$. In this paper we focus on the 'giant' glitches, i.e. glitches with fractional increases in the spin rate of the order of $\\frac{\\Delta\\nu}{\\nu}\\approx 10^{-6}$, that are observed in a sub class of pulsars including the Vela. We show that giant glitches can be modelled with a two-fluid hydrodynamical approach. The model is based on the formalism for superfluid neutron stars of Andersson and Comer (2006) and on the realistic pinning forces of Grill and Pizzochero (2011). We show that all stages of Vela glitches, from the rise to the post-glitch relaxation, can be reproduced with a set of physically reasonable parameters and that the sizes and waiting times between giant glitches in other pulsars are also consistent with our model. ",
        "introduction": "The timing of radio pulsars provide us with one of the most stable clocks in the universe. Many pulsars exhibit, however, sudden increases in the spin rate, known as 'glitches'. To date there have been glitches reported in more than 100 pulsars\\footnote{http://www.jb.man.ac.uk/pulsar/glitches/gTable.html} \\citep{Espinoza}, with fractional jumps in the spin rate of $\\frac{\\Delta\\nu}{\\nu}\\approx 10^{-11}$ to $10^{-5}$.  There is still no clear consensus on the origin of these phenomena, but it is thought that glitches (at least the giant glitches that are observed in somewhat older, colder pulsars such as the Vela) are due to angular momentum being stored in a superfluid component of the star that is temporarily decoupled from the charged component to which the electromagnetic emission is anchored. When the two components recouple there is a sudden transfer of angular momentum to the crust, which gives rise to the observed spinup \\citep{Anderson}. A superfluid rotates by forming a quantised array of vortices, the distribution of which determines the rotational profile of the star. The main idea at the base of most glitch models is that vortices can {\\it{pin}} to the crustal lattice \\citep{Anderson,Alpar77,Pines80,Alp81,Anderson82}, allowing for a lag to build up between the superfluid and the charged component.  Following the seminal work of \\citet{BPP69} several models have been developed to explain the dynamical evolution of the pinned superfluid coupled to the crust, mainly to with the intent of explaining the observed timescales of days to months in the post-glitch recovery of the Vela pulsar. Essentially one can distinguish two classes of models, those that assume that the relaxation is due to the weak coupling between the superfluid and the crust due to the interaction between free vortices and the coulomb lattice of nuclei \\citep{Jones2, JonesGL, JonesGL2} and those which rely on the creep model based on the seminal work of  \\citet{Alpar1} and later adaptations \\citep{ALS,LEB}. In this case the assumption is that as a lag develops between the crust and the superfluid, vortices {\\it creep} through the crust at a rate that is highly temperature dependent, gradually transferring angular momentum and leading to a steady state spindown. However, when the lag approaches a critical value, it is possible for an instability to trigger a catastrophic unpinning event associated with a sudden transfer of angular momentum to the crust, after which vortices repin and a new equilibrium is reached. These models work remarkably well in explaining the post-glitch relaxation, but struggle to explain the exact nature of the perturbation which triggers it. Such an event may be triggered by large temperature perturbations \\cite{LE96}, caused for example by starquakes \\cite{BP71,Cheng92}, the interactions of the proton vortices and the crustal magnetic field \\citep{SC99}  or, in the presence of strong crustal pinning, a two-stream instability in the superfluid flow may develop, leading to a sudden transfer of angular momentum \\citep{kglitch} Recently, however,  \\citet{NewM} have shown that glitches can be triggered by unpinning avalanches of the vortex array, leading to very promising results for the overall glitch size and waiting time distributions (see also \\citet{Melatos1, Melatos2}).   An entirely different mechanism was proposed by \\citet{Rud91}, who suggested that  vortices pinned to the crust stress it to the point of fracture and then move outwards with the matter they are pinned to. Several other ``starquake'' models have been proposed \\citep{EL2000,FLE2000} but they are all based on the main assumption that a starquake triggers the outward motion of vortices pinned to the matter and the local rearrangement of the crust then causes an increase of the angle between the rotation axis and the magnetic axis, leading to an increased electromagnetic braking torque on the star. Note that starquakes had been invoked as a possible explanation for pulsar glitches very soon after the first observations, but had been ruled out as they could not account for the giant glitches of the Vela \\citep{BP71,Rud75}.  This mechanism may, however, play a role in other systems and in fact there is some evidence that smaller glitches in younger and hotter pulsars such as the Crab may be linked to starquakes \\citep{Mid}. Furthermore it has been suggested that the glitch may not be exclusively linked to vortex dynamics in the crust, but that the interaction between the vortices and the superconducting flux tubes in the core may contribute to the event \\citep{Rud98,trevnew}.   One of the main difficulties in constructing a glitch model is, however, the relative lack of realistic calculations for the pinning forces. Several calculations exist of the pinning {\\it energy} and of the force per pinning site \\cite{Alpar77, Alpar1,EB88,LE91, PV,PBJ98, DP03, DP04, DP06,Avo} but such calculations of the pinning force neglect the finite length of the vortices and rely on simplified geometries. Recently, however, \\citet{Linkn} has analysed the motion of a vortex in a random potential and, even more crucially, \\citet{pinning1} have obtained realistic estimates of the pinning force {\\it per unit length} of the vortex, the quantity that can be directly compared to the Magnus force which tends to push the vortices out and force the superfluid into corotation with the crust.    \\citet{Pierre1} has recently shown that a simple analytic model, based on the catastrophic unpinning paradigm and the realistic pinning forces calculated by \\citet{pinning1} (see also \\citet{pinning1phd}), can describe glitches in the 'Vela' like pulsars, that are typically older, exhibit larger glitches ($\\frac{\\Delta\\nu}{\\nu}\\approx 10^{-7}$ to $10^{-6}$) and always show a decrease in the frequency derivative after the glitch \\citep{Espinoza}. Essentially the assumption is that as the star spins down the vorticity is accumulated in the strong pinning region and then released once the resulting Magnus force exceeds the maximum of the pinning force.   In order to make quantitative predictions for the observed jump in frequency and subsequent relaxation of the crust it is, however, also necessary to describe the evolution of the various fluid components that form the NS. In order to do this it is convenient to follow a hydrodynamical approach that does not deal with vortex motion directly, but rather follows the evolution of the separate fluid components. In particular we shall make use of the multifluid formalism developed by \\citet{NilsandGreg} and recently applied, albeit in a much simpler formulation, to the study of pulsar glitches by \\citet{glitch}. In the following we shall thus go beyond the analytic model of \\citet{Pierre1} and show how the realistic pinning results of \\citet{pinning1} can be incorporated into this formalism and used to accurately reproduce, for a reasonable range of parameters, the observed properties of pulsar glitches. Furthermore, as we will discuss in the following sections, we will assume that the relaxation is entirely due to the mutual friction between the superfluid component and the crust. In a sense we are thus taking the view of \\citet{JonesGL} that the relaxation is due to the dynamics of vortices close to corotation with the superfluid, rather than to vortex 'creep'. Note, however, that the creep model has been shown to be consistent with the observed relaxation timescales of the Vela pulsar in the linear regime \\citep{Alp89},  and we can thus qualitatively account for it by rescaling the mutual friction parameters, as will be discussed in the following. Finally let us also point out that an effect that is neglected in this preliminary model is that of Ekman pumping at the crust-core interface, which has been shown, together with shear viscosity, to be relevant in explaining the post-glitch relaxation times \\citep{vE1} and should be included in future developments.   Let us stress that such detailed theoretical work is crucial at this time, as the new low frequency radio telescope LOFAR has just come online and begun observations of pulsars and fast radio transients \\citep{LOFAR}. As development continues LOFAR, and in the future the Square Kilometer Array (SKA), are likely to not only provide a wealth of data on both known and new glitching pulsars, but also resolve (or at least tightly constrain) the glitch rise time, thus allowing us to test our theoretical understanding of the glitch mechanism \\citep{SKA}. Furthermore a glitch may trigger a gravitational wave (GW) burst  which could be detectable by next generation GW observatories \\citep{glitch,vE2}. By accurately determining the glitch time LOFAR and the SKA could be used to trigger a directed GW search.  Finally the constructions of long-term hydrodynamical simulations of NS interior could also be extended to model not only large pulsar glitches, but more generally pulsar timing noise, in order to understand the role that the interior dynamics of the star plays compared to magnetospheric processes, which have been shown to be at work in some systems \\citep{Lyne}. This is an issue that is of great importance for the current efforts to detects GWs with pulsar timing arrays \\citep{Hobbs}. ",
        "conclusions": "We have presented a hydrodynamical two fluid model of pulsar glitches that can consistently model all phases of the glitch itself. Our model can be successfully applied to the giant glitches of the Vela pulsar, for which we can reproduce the approximate waiting time between glitches, the size of the glitch and the short term post-glitch relaxation. The main assumption is that a giant glitch will occur once the system exceeds the maximum lag that the pinning force in the crust can sustain. This naturally gives rise to a waiting time between glitches that depends on the pulsar spindown rate (i.e. it is the time it takes the crust to spin down by the required amount) and to a maximum size for the glitch. Both these quantities depend only weakly on the mass and radius of the star \\citep{pinning1, pinning1phd, Stefano} and our model can, in fact, reproduce these features successfully also for the general population of \"giant glitchers\", i.e. the pulsars for which giant glitches such as those of the Vela have been observed \\citep{Espinoza}.   In our model the coupling between the charged component (which we assume to be rigidly rotating) and the superfluid neutrons is given by the vortex-mediated mutual friction.  Our results suggest that the mutual friction will be weak in the crust, possibly due to the fact that not all vortices are free, but rather that the strong pinning force gives rise to a situation in which most vortices are pinned and only a small fraction can 'creep' outwards. Only once the maximum unpinning lag is exceeded can the vortices move out freely; a process which can excite Kelvin oscillations and give rise to a strong drag and recoupling of the two components on a very short timescale, i.e. a glitch. The short term post-glitch relaxation of the Vela, on the other hand, suggests that the magnitude of the drag in the core of the NS is consistent with theoretical expectations for electron scattering of magnetised vortex cores. Our model does not support the notion that, at least on short timescales, a significant number of vortices is pinned in the core (as could, for example, be the case if one has a type II superconductor and vortices cannot cross fluxtubes, effectively decoupling the core and the crust). A detailed analysis of the case in which the core consists of a type II superconductor will be a focus of future work in order to obtain more quantitative results and constraints on NS interior physics.  Some vortices that cross the core  may however be weakly pinned to the crust, and vortex repinning and creep (also in the core) may play a role on the longer timescales associated with the recovery.  Another effect which will have an impact on the post-glitch recovery is the Ekman flow at the crust-core interface. This effect has been shown to be important in fitting the post glitch recovery of the Vela and Crab pulsars by \\citet{vE1} and future adaptations of our model should relax the rigid rotation assumption for the charged component and include the effect of Ekman pumping.  Further developments should also include more realistic models for the drag parameters in the star, as the density dependence of the coupling strength clearly has an impact on the amount of angular momentum that can be exchanged on different timescales. Truly quantitative results could then be obtained with the use of realistic equations of state together with consistent estimates of the pinning force, such as those of \\citep{pinning1} and \\citep{pinning1phd}.   Note that we have assumed that a giant glitch only occurs when the maximum critical lag is reached. If unpinning could be triggered earlier, this could generate smaller glitches. In fact cellular automaton models have shown that the waiting time and size distributions of pulsar glitches can be successfully explained by vortex avalanche dynamics, related to random unpinning events \\citep{Melatos1, Melatos2,NewM}. It would thus be of great interest to use our long-term hydrodynamical models, with realistic pinning forces, as a background for such cellular automaton models that model the short-term vortex dynamics. Such a model could then also be extended to model not only large pulsar glitches, but more generally pulsar timing noise, an issue that is of great importance for the current efforts to detects GWs with pulsar timing arrays \\citep{Hobbs}.   Finally, the next generation of radio telescopes, such as LOFAR and the SKA, is likely to provide much more precise timing data for radio pulsars and is likely to set much more stringent constraints on the glitch rise time and short term relaxation, thus allowing us to test our models and probe the coupling between the interior superfluid and the crust of the NS with unprecedented precision."
    },
    "1107/1107.2226_arXiv.txt": {
        "abstract": "{We present VLT/UVES spectroscopy of the quasar Q0841+129, whose spectrum shows a proximate damped \\lya\\ (PDLA) absorber at $z$=\\zdla\\ and a proximate sub-DLA at $z$=\\zsub, both lying close in redshift to the QSO itself at $z_{\\rm em}$=2.49510$\\pm$0.00003. % This fortuitous arrangement, with the sub-DLA acting as a filter that hardens the QSO's ionizing radiation field, allows us to model the ionization level in the foreground PDLA, and provides an interesting case-study on the origin of the high-ion absorption lines \\sif, \\cf, and \\os\\ in DLAs. The high ions in the PDLA show at least five components spanning a total velocity extent of $\\approx$160\\kms, whereas the low ions exist predominantly in a single component spanning just 30\\kms. We examine various models for the origin of the high ions. Both photoionization and turbulent mixing layer models are fairly successful at reproducing the observed ionic ratios after correcting for the non-solar relative abundance pattern, though neither model can explain all five components. We show that the turbulent mixing layer model, in which the high ions trace the interfaces between the cool PDLA gas and a hotter phase of shock-heated plasma, can explain the average high-ion ratios measured in a larger sample of 12 DLAs.} ",
        "introduction": "In the process of galaxy assembly and evolution, interstellar and circumgalactic gas clouds play a key role, representing both the fuel for future star formation and the metal-enriched by-products of past star formation. Absorption-line spectroscopy is a powerful way to study these gas reservoirs, to characterize their kinematic and chemical properties, and to investigate their link to star formation. Of particular interest for probing interstellar gas are the damped \\lya\\ systems (DLAs), the highest column density absorbers defined by log\\,$N$(\\hi)$>$20.3 \\citep{Wo86}. Their high neutral gas content ensures they select over-dense regions, including galaxies and their halos. DLAs contain the majority of the neutral gas in the Universe in the range $z$=0--5 \\citep{Wo05}, and serve as luminosity-independent indicators of cosmic chemical enrichment.  Despite being known for their neutrality, DLAs frequently show high-ion absorption overlapping in velocity space with the neutral lines yet tracing a separate phase (or phases) of gas. The first high ions detected in DLAs were \\cf\\ and \\sif\\ \\citep{Wo94, Lu95, Lu96, WP00a, WP00b}, which almost always show a different (though overlapping) component structure than the low ions. We have recently been surveying high-ion absorption in DLAs at $z$=2--4 at high resolution using the VLT/UVES spectrograph, finding the high ions are ubiquitous. In data with S/N$\\ga$30 per pixel, we report detection rates of $>$34\\% for \\os\\ \\citep{Fo07a}, 100\\% for \\cf\\ and \\sif\\ \\citep{Fo07b}, and 13\\% for \\nf\\ \\citep{Fo09}. High-ion absorption is also seen in a composite SDSS spectrum of several hundred DLAs \\citep{Ra10}. In \\citet{Fo07a} we favored the interpretation that the \\os\\ absorption components in DLAs trace warm-hot ($T\\!\\ga\\!10^5$\\,K) collisionally-ionized plasma, rather than cool ($T\\!\\sim\\!10^4$~K) photoionized plasma, based on two observational findings: first, a tendency for the individual \\os\\ components in DLAs to be broader than the \\cf\\ components (suggesting higher temperatures in the \\os\\ phase), and second, the failure of photoionization models to reproduce the \\os\\ column densities measured in a particular DLA studied in detail (the \\zabs=2.076 DLA toward \\object{Q2206--199}) without recourse to unreasonably large path lengths. \\citet{Le08} investigated the \\zabs=2.377 DLA toward \\object{J1211+0422}, and similarly found that photoionization cannot reproduce the observed \\os\\ column, although they also found that single-phase collisional ionization models (equilibrium or non-equilibrium) cannot explain the high-ion column densities in that system.  The origin of \\os\\ absorption in DLAs demands further investigation, since the observational characterization of warm-hot gas in galaxy halos at redshifts 2--3 would provide useful constraints on models of galaxy formation and evolution. Unfortunately, observations of \\os\\ absorption at these redshifts are not straightforward, since the \\os\\ resonance doublet ($\\lambda\\lambda$1031,1037) lies in the \\lya\\ forest and so often suffers from severe contamination from foreground \\hi\\ absorbers \\citep[e.g.][]{BT96, Fr10}. With a large enough DLA sample, one occasionally comes across a case where the \\os\\ doublet lines are only partially (rather than completely) blended, and where information on portions of the \\os\\ line profiles can be recovered \\citep{Fo07c, Le08, Si02}. Such a case (the DLA at \\zabs=\\zdla\\ toward \\qa) is discussed in this paper. This system has another advantage: although it lies at only 1\\,630\\kms\\ from the background % QSO (classifying it as a proximate DLA or PDLA), it is shielded from the QSO's ionizing radiation by a second absorber, a sub-DLA at \\zabs=\\zsub, very close in redshift to the QSO at \\zem=\\zqso. The \\hi\\ column in the sub-DLA, log\\,$N$(\\hi)=19.00$\\pm$0.15, is high enough that the system is optically thick at the Lyman Limit, but not at high photon energies ($E\\!\\ga\\!75$\\,eV, i.e. $\\lambda\\!\\la\\!165$\\,\\AA). Therefore the sub-DLA acts to harden the ionizing radiation field incident on the PDLA.  PDLAs are defined as those DLAs lying within 5\\,000\\kms\\ (or sometimes 3\\,000\\kms) of the redshift of the background QSO \\citep{El02, El10, El11, Ru06, Rx07, Pr08}. They are frequently excluded from studies of intervening DLAs, due to their potentially unusual ionization conditions and environmental effects (such as galaxy clustering near quasars). However, the fact that the (enhanced) ionizing radiation field in PDLAs is calculable can be turned to our advantage, and allows them to be used as case-studies for the origin of high-ionization plasma in DLAs. That is the approach followed in this paper. ",
        "conclusions": "In order to explore the processes that generate high-ion plasma in high-redshift galaxy halos, we have presented new VLT/UVES high-resolution observations of the PDLA at \\zabs=\\zdla\\ toward \\qa, focusing on the origin of the high-ion absorption. This system lies at 1\\,630\\kms\\ from the % background QSO, has log\\,$N$(\\hi)=20.80$\\pm$0.10, a metallicity [Zn/H]=$-$1.60$\\pm$0.10, a low-ion velocity width $\\Delta v_{90}$=30~\\kms, and a hydrogen ionization fraction in the low-ion phase of 0.1--0.5 as determined from the \\aro/\\siw, \\aro/\\sw, \\nw/\\none\\ and \\fet/\\few\\ line ratios. The high ions in the PDLA show a multi-component structure; at least five \\cf\\ components are seen spanning a total velocity extent of $\\approx$160\\kms\\ with $b$-values ranging from 8.8$\\pm$1.8\\kms\\ to 27.3$\\pm$3.1\\kms. \\sif\\ is detected in three of these components; \\os\\ is detected in two, and blended in the others.  The PDLA is (partly) shielded from the QSO's ionizing radiation by an intrinsic sub-DLA at \\zabs=\\zsub, very close to the QSO emission-line redshift of \\zem=\\zqso. This filtering arrangement hardens the ionizing radiation field incident on the PDLA, allowing us to test whether photoionization (PI) is a viable origin mechanism. We find that a single-phase PI model at log\\,$n_{\\rm H}\\!\\approx\\!-2.9$ can successfully explain one of the two components showing \\os\\ (C1), so long as [C/O]=$-$0.5, but no single-phase PI model is found for the other component showing \\os\\ (C2), since the \\os\\ in this component is a factor of four stronger than the model predicts. The \\cf\\ and \\sif\\ profiles in C2, C3, and C4, are consistent with PI at log\\,$n_{\\rm H}\\!\\approx\\!-2$ if [C/Si]=$-$0.5. Single-phase collisional ionization models (equilibrium or non-equilibrium) are unable to reproduce the high-ion observations.  An alternative model, where the high ions trace the turbulent mixing layers between the warm DLA gas and a hot exterior medium \\citep{KS10}, can successfully reproduce the ionic ratios in four of the five components (C2, C3, C4, and C5). This model can also explain the average high-ion ratios \\cf/\\os, \\nf/\\os, and \\cf/\\sif\\ observed in a larger sample of 12 DLAs. Turbulent mixing layers are often invoked (along with conductive interfaces) to explain high-ion observations in the halo of the Milky Way (itself a DLA). If correct, this model provides indirect evidence for the existence of a hot phase of coronal plasma in and around DLA galaxies. Further observations of high-ion absorption in individual intervening DLAs are needed to test this model."
    },
    "1107/1107.5630_arXiv.txt": {
        "abstract": "We report a spectroscopic analysis of an EIT wave event that occurred in active region 11081 on 2010 June 12 and was associated with an M2.0 class flare. The wave propagated near circularly. The south-eastern part of the wave front passed over an upflow region nearby a magnetic bipole. Using EIS raster observations for this region, we studied the properties of plasma dynamics in the wave front, as well as the interaction between the wave and the upflow region. We found a weak blueshift for the \\ion{Fe}{12} $\\lambda$195.12 and \\ion{Fe}{13} $\\lambda$202.04 lines in the wave front. The local velocity along the solar surface, which is deduced from the line of sight velocity in the wave front and the projection effect, is much lower than the typical propagation speed of the wave. A more interesting finding is that the upflow and non-thermal velocities in the upflow region are suddenly diminished after the transit of the wave front. This implies a significant change of magnetic field orientation when the wave passed. As the lines in the upflow region are redirected, the velocity along the line of sight is diminished as a result. We suggest that this scenario is more in accordance with what was proposed in the field-line stretching model of EIT waves. ",
        "introduction": "Diffuse coronal waves were first observed by \\citet{mos97} and \\citet{tho98} with the EUV Imaging Telescope (EIT; \\citealt{dela95}) aboard the {\\it Solar and Heliospheric Observatory} ({\\it SOHO}), and were commonly known as EIT waves. They are best seen in the running difference images as propagating bright fronts with a speed of a few hundreds of km s$^{-1}$, followed by an expanding dimming region \\citep{tho98}. EIT waves can be observed at several wavelengths, such as 171 \\AA, 195 \\AA, 284 \\AA, and 304 \\AA~\\citep{wil99,zhu04,lon08}. Many observational properties of EIT waves were presented by \\citet{del99}, \\citet{kla00}, and \\citet{tho09}. Recent reviews on this topic can be found in \\citet{wil09}, \\citet{war10}, and \\citet{gall10}.  It is natural that EIT waves are accompanied with some other solar active events such as coronal mass ejections (CMEs) and flares. Statistical and case studies have shown that EIT waves are closely related with CMEs rather than solar flares \\citep{bie02,oka04,chen06}. In particular, the flares associated with many EIT wave events were very weak \\citep{cliv05,ver08,ma09,att09}. In spite of the coincidence of EIT waves and CMEs, their spatial relationship is still being debated. \\citet{vrs06} reported a wave behind the CME flank, and \\citet{pat09} claimed that the EIT wave front is outside the CME frontal loop, whereas \\citet{chen09} and \\citet{dai10} found that the EIT wave front is cospatial with the white-light frontal loop of CMEs. In addition, \\citet{war10} found one event in which a wave front overtakes the CME flank.  Based on the observed properties of EIT waves, there are two main interpretations. One is the wave model, which suggests that EIT waves are fast-mode magnetohydrodynamic (MHD) waves or shocks \\citep{wang00,wu01,war01,war04,pom08,vrs08}. This model can explain some of the characteristics of the wave front and was supported by a number of observations \\citep{war05,lon08,ver08,pat09,gopa09,kie09}. The other interpretation is related to CME expansion. \\citet{del00} suggested that the bright front may result from the interaction between CME-induced expansion of magnetic field lines and surrounding field lines. \\citet{chen02,chen05} proposed a field-line stretching model, which predicts that there should exist a fast moving coronal shock ahead of the slow EIT wave, which were recently confirmed by \\citet{chen11}. The non-wave model can also explain many of the properties of the wave and was supported by \\citet{har03} and \\citet{zhu09}. In addition, there are still other models such as slow-mode MHD waves \\citep{wang09}, spherical current shell \\citep{del08}, successive magnetic reconnections \\citep{att07a,att07b}, and soliton \\citep{wil07}. A concept of a coupled coronal wave was proposed by \\citet{zhu04} and \\citet{coh09}.  In the last decade, the main approach to studying EIT waves is by ultra-violet imaging observations. The main limitation comes from the cadence and spatial resolution of the instruments, such as EIT on board the {\\it SOHO} and Extreme UltraViolet Imager (EUVI) on board the {\\it Solar TErrestrial RElations Observatory} ({\\it STEREO}; \\citealt{how08}). Atmospheric Imaging Assembly (AIA; \\citealt{tit06}) aboard the {\\it Solar Dynamics Observatory} ({\\it SDO}) has high spatio-temporal resolutions and signal-to-noise ratio, enabling us to study EIT waves in unprecedented detail. \\citet{liu10} reported the first {\\it SDO}/AIA observations of EIT wave. They found one diffuse pulse and multiple sharp fronts and suggested a hybrid interpretation, combining both wave and non-wave models. In addition to the imaging observations, spectroscopic observations are also very important, because they provide additional information on plasma dynamics during the wave propagation and aid  clarification of the physical nature of EIT waves. \\citet{har03} did the first spectroscopic analysis of EIT waves using Coronal Diagnostic Spectrometer (CDS; \\citealt{harrison95}) on board the {\\it SOHO} spacecraft. They measured Doppler velocities in the wave front and the following dimming. They found an absence of Doppler velocity in the wave front but an ejection of cold material after the wave front passed. Using the EUV Imaging Spectrometer (EIS; \\citealt{cul07}) on {\\it Hinode}, \\citet{asa08} studied a fast mode MHD shock wave visible in soft X-rays. Unfortunately, some of the EIT images suffered from scattered light in the telescope. Therefore, the wave front in EIT data was unclear. More recently, \\citet{chen10} confirmed the studies of \\citet{har01,har03}, and further found an enhanced line broadening at the outer edge of the dimming region, which could be well explained by the field-line stretching model of \\citet{chen02,chen05}.  In this paper,we present a case study of {\\it Hinode}/EIS observation of an EIT wave event. We successfully obtained the temporal evolution of the line intensity, line width, and Doppler velocity for two iron lines in a sliced region overlapping a small upflow region. Hence, we are able to reveal the interaction between the EIT wave and the coronal upflow region spectroscopically for the first time. We describe the observations and data analysis in \\S \\ref{obda}. Our results are shown in \\S \\ref{result}, followed by some discussions on the results in \\S \\ref{discu}. ",
        "conclusions": "\\label{discu}  We presented an EIT wave event, which was captured by {\\it Hinode}/EIS raster. We studied the spectroscopic properties of the coronal plasma during the wave transit. The most interesting results are summarized below.  \\begin{enumerate} \\item{The wave front propagated near circularly, with its south-eastern part observed by EIS. The line intensity and line width showed no change when the wave front passed. However, an enhanced blueshift of the line center, albeit weak, was observed at this time.} \\item{The wave front passed over an upflow region nearby a magnetic bipole. The shape of the wave front was only slightly distorted. For both of the iron lines studied, a sudden and clear diminishing of the upflow and non-thermal velocities in the upflow region was observed when the wave passed.} \\end{enumerate}  In general, during the EIT wave front passage, one may expect to observe a line intensity enhancement, a common feature observed in coronal EUV images.  In this event, however, the variation of line intensity during the wave propagation was within the fitting error. For the raster observations of the wave, the line profile at a certain point is a composite of the contribution from the background corona during the whole exposure time (50 s for this event) and the contribution from the wave front. The wave quickly passes over this point and in fact it contributes little to the line profile. Therefore, it is hard to observe similar signatures from spectroscopic and image observations. It was previously observed that EIT waves and the accompanied dimmings exhibited a very high speed, so that its contribution to the line profile can result in a strong Doppler component \\citep{har03,asa08}. In this case, some significant features of the wave or wave-perturbed plasma can be observed from spectroscopic data.  Fortunately, we found a sudden upward motion of the coronal plasma when the wave front passed the EIS FOV. The upward velocity ($v_{up}$) was $\\sim$~4 km s$^{-1}$. If we assume that the velocity along the line of sight was purely from the projection of a tangential velocity (i.e., velocity along the solar surface), then we can estimate the local tangential velocity to be $v_{up}$/$\\cos\\theta$, where $\\theta$, the projection angle, is $\\sim$~30$\\degr$. Thus, the local tangential velocity was $\\sim$~10 km s$^{-1}$. For this event, the average propagation speed of the wave between 1:00:54 and 1:05:30 UT is $\\sim$~358 km s$^{-1}$. This result implies that the speed of local plasma motion was far less than the propagation speed of the wave. This feature is compatible with both the wave models and non-wave models.  The upflows at the edge of the active regions have been reported by \\citet{sakao07}, \\citet{hara08}, \\citet{har08}, and \\citet{dosc08}. They were found to be cospatial with open magnetic field lines \\citep{sakao07,har08}. The upflow region observed in our event, as described in Section \\ref{result_upflow}, had some similar behaviors. However, its spatial extension was much smaller. Base on the fact that the upflow we observed existed in a low intensity region, it is probable that this region was a small coronal hole. Regardless whatever the actual source of the upflow was, the magnetic field lines in this region were very likely to be open. However, several difficulties make it hard to get a valid result on the magnetic field structure for this event. First, both the active region where the EIT wave originated and the magnetic bipole region where upflows existed were far from the disk center. A correction for the curved surface and the projection of magnetic fields is required. Second, although extrapolation to active region magnetic fields has been intensively studied and relatively easy to apply to various cases, extrapolation to small magnetic structures is a difficult task. Unfortunately, the region with the magnetic bipole and the upflows, which draws our main attention, is such an example. Our experiment showed that it is possible to get some preliminary results, though probably inaccurate, for the active region. However, no valid result is returned for the magnetic bipole or the intermediate area between it and the active region. Although there is no available information on the magnetic fields, we conjecture that the upflow region surrounding the magnetic bipole core may also be related with open fields or large-scale close fields, based on the reasons mentioned above.  The ratio of magnetic pressure to gas pressure is high in the corona; thus, the direction of coronal plasma motion is dominated by magnetic field lines. The sudden variation of the line of sight velocity described in Section \\ref{result_upflow} implies the change of direction of magnetic field lines. \\citet{mcin07} reported a disappearance and reappearance of ``moss\" around an active region during the evolution of a CME event. They explained the change of the ``moss\" as a proxy of the changing coronal magnetic field topology behind the CME front. Although their event is different from ours in some aspects, the common key feature is that a change of magnetic field orientation associated with large scale coronal disturbances (CMEs or EIT waves) can influence the coronal plasma dynamics. In the field-line stretching model proposed by \\citet{chen02,chen05}, EIT waves are apparently-propagating wave fronts formed by successive stretching/expansion of field lines, which is initiated by the erupting flux rope. It is expected that when the wave passed over the upflow region, the local open (or large-scale) field lines would be pushed aside and be redirected toward another direction by the stretched field lines related with the wave front. Therefore, the upflows changed their direction to that deviated more from the line of sight. The line of sight component velocity were diminished apparently as a result.  \\citet{hara08} and \\citet{dosc08} reported a positive correlation between the non-thermal velocity and upflow velocity from the line profile fitting, implying the multiplicity of actual upflow velocities. As shown in Figure \\ref{velocity}, before the wave front passed, the upflow velocity and non-thermal velocity in the region we observed also shows a positive correlation, though somewhat weaker than what was reported by \\citet{hara08} and \\citet{dosc08}. However, such a correlation does not exist after the wave front passed, when the upflow and non-thermal velocities were reduced to smaller magnitudes as mentioned in Section \\ref{result_upflow}. Considering the difference in the correlation result, we think that the correlation observed in this upflow region before the wave front passed is of physical significance, especially for the \\ion{Fe}{13} line. Since one of the main causes of the non-thermal velocity is the multiplicity of line of sight velocities, the sudden diminishing of the non-thermal velocity could therefore be explained. Note that other mechanisms were also suggested to explain the variation of non-thermal velocity in coronal dimmings \\citep{mcin09}. Here, we favor the scenario of multiple velocities existing in the spatially unresolved area to interpret the variation of the non-thermal velocity.  Note that an MHD wave impacting open (or large-scale) field lines can also push them aside and redirect them to some extent. However, the field lines would oscillate periodically. If so, we expect to observe some oscillating patterns in the spectroscopic data. Unfortunately, no oscillation is found in the results shown above. It may be true that the sensitivity and the cadence of the instruments are not high enough to resolve the oscillation well. If this is the case, the diminishing of the Doppler velocity and non-thermal velocity would be gradual. However, Figure \\ref{eis} shows a sudden diminishing of the velocities for both the iron lines. Therefore, while the wave models cannot be excluded, the non-wave model is more favored here in explaining the observational data.  In summary, a possible scenario for the EIT wave event analyzed in this paper could be as follows: the EIT wave encountered an open (or large-scale) magnetic field line, along which multi-component upflows existed. The independent magnetic system of the upflow region was too narrow to stop the wave from propagating. However, the open (or large-scale) field lines were pushed aside or distorted by the wave, resulting in an abrupt diminishing of upward line of sight velocity and non-thermal velocity as observed in spectral lines. We suggest that this scenario could be more in accordance with what proposed by \\citet{chen02,chen05} for EIT waves, in which the formation of EIT wave fronts and their behaviors are essentially correlated with the stretching of magnetic field lines during CMEs."
    },
    "1107/1107.0824_arXiv.txt": {
        "abstract": "We present new GMRT observations of HDF\\,130, an inverse-Compton (IC) ghost of a giant radio source that is no longer being powered by jets. We compare the properties of HDF\\,130 with the new and important constraint of the upper limit of the radio flux density at $240$~MHz to an analytic model. We learn what values of physical parameters in the model for the dynamics and evolution of the radio luminosity and X-ray luminosity (due to IC scattering of the cosmic microwave background (CMB)) of a Fanaroff-Riley II (FR~II) source are able to describe a source with features (lobe length, axial ratio, X-ray luminosity, photon index and upper limit of radio luminosity) similar to the observations. HDF\\,130 is found to agree with the interpretation that it is an IC ghost of a powerful double-lobed radio source, and we are observing it at least a few Myr after jet activity (which lasted $5$--$100$~Myr) has ceased. The minimum Lorentz factor of injected particles into the lobes from the hotspot is preferred to be $\\gamma\\sim10^3$ for the model to describe the observed quantities well, assuming that the magnetic energy density, electron energy density, and lobe pressure at time of injection into the lobe are linked by constant factors according to a minimum energy argument, so that the minimum Lorentz factor is constrained by the lobe pressure. We also apply the model to match the features of 6C\\,0905+3955, a classical double FR~II galaxy thought to have a low-energy cutoff of $\\gamma\\sim10^4$ in the hotspot due to a lack of hotspot inverse-Compton X-ray emission. The models suggest that the low-energy cutoff in the hotspots of 6C\\,0905+3955 is $\\gamma\\gtrsim 10^3$, just slightly above the particles required for X-ray emission. ",
        "introduction": "In this work we establish whether a consistent interpretation can be found for the currently observed properties of the double-lobed sources HDF\\,130 and 6C\\,0905+39, the former of which is thought to no longer have current jet activity. We identify at which stage in their life cycles HDF\\,130 and 6C\\,0905+3955 could be by comparing their observable features (lobe length, axial ratio, X-ray luminosity, photon index, radio flux density limit) to an analytic model for the dynamics and evolution of X-ray and radio emission of an active FR~II object whose jets switch off after a time $t_{\\rm j}$ developed in \\citet{2011MNRAS.413.1107M}. The models also help suggest what the physical parameters of the sources such as jet energy, jet lifetime, ambient density parameters, and injection spectrum parameters may be, some of which are difficult to determine from observation alone.  {\\it Hubble Deep Field}~(HDF)~130 is an extended X-ray source observed in the {\\it Chandra Deep Field}-North X-ray image \\citep{2003AJ....126..539A}. All six of the extended  X-ray sources found in the $1$~Ms exposure were attributed to clusters and groups by \\cite{2002AJ....123.1163B}. However, HDF\\,130 has since been realised to be a double-lobed structure with extended X-ray emission due to IC scattering of the CMB \\citep{2009MNRAS.395L..67F}. While the jet is turned on in a powerful radio source, $\\gamma\\sim10^4$ electrons (assuming typical magnetic field strengths) required to generate GHz synchrotron radiation in the radio band lose their energy due to radiative losses much more quickly than the $\\gamma\\sim10^3$ electrons responsible for upscattering the CMB photons; these losses are compounded by the expansion of the plasma. Thus after jet activity ceases, the IC X-ray emission lasts longer than radio emission and the source will appear as an IC ghost of a radio lobe for some period of time. On a morphological basis, and also given the non-thermal nature of the spectrum of the extended X-ray emission from HDF\\,130 \\citep{2009MNRAS.395L..67F}, the extended source is most likely the lobes of a formerly powerful Fanaroff-Riley II (FR~II) \\cite{1974MNRAS.167P..31F} galaxy.  HDF\\,130 is approximately $690$~kpc across as determined by \\cite{2009MNRAS.395L..67F}, and such an extent is not exceptional (e.g. \\citealt{2008MNRAS.390..595M}). The extended emission has a steep photon index of $\\Gamma=2.65$, which could indicate significant synchrotron cooling. HDF\\,130 is about half as bright at X-ray wavelengths as the giant powerful radio galaxy 6C\\,0905+3955 \\citep{2008MNRAS.386.1774E}, which has a spectral index of $\\Gamma=1.61$, suggesting that HDF\\,130 may be viewed at a later stage in its life cycle than 6C\\,0905+3955.  6C\\,0905+3955 is a powerful FR~II galaxy, approximately $945$~kpc in diameter \\citep{2006ApJ...644L..13B}. The source characteristics mandate a low-energy cutoff of freshly injected particles in the hotspot above $\\gamma\\sim10^3$ due to the absence of X-ray emission from the hotspot (but not the lobes) \\citep{2006ApJ...644L..13B,2008MNRAS.386.1774E}. The lobes, containing older plasma, do have $\\gamma\\sim10^3$ particles required for observing IC scattering on the CMB in the X-ray. The higher energy particles injected from the hotspot into the lobes undergo energy loss and hence result in the presence of plentiful $\\gamma\\sim10^3$ particles in the lobes. Extended IC X-ray emission has been observed in other sources as well, such as 3C\\,294 ($z=1.786$) \\citep{2003MNRAS.341..729F}, 4C\\,23.56 ($z=2.48$) \\citep{2007MNRAS.376..151J} and 4C41.17 ($z=3.8$) \\citep{2003ApJ...596..105S}. The CMB energy density is proportional to $(1+z)^4$, cancelling the dimming due to distance, and thus extended X-ray emission may be observable at both low and high redshifts \\citep{1969Natur.221..924F}. ",
        "conclusions": "\\label{sec:Discussion}  We matched the observational features of HDF\\,130 and 6C\\,0905+3955 to analytic models of the evolution of the lobe length, axial ratio, photon index, and X-ray and radio luminosities of an FR~II object. {\\it Only periods of time when the source is no longer active in the models are congruent to the observations of HDF\\,130, supporting the idea that HDF\\,130 is an IC ghost of a giant radio source.} The models suggest HDF\\,130 experienced jet activity for a period of around $5$--$100$~Myr, and that we are viewing the object at least a few Myr after the jets have turned off, which is why the source is not bright in the radio. 6C\\,0905+3955 is inferred to have had an active jet for at least $20$~Myr and may have a slightly higher intrinsic jet power than HDF\\,130.  Predicted radio luminosities at lower frequencies ($151$ and $74$ MHz) for the models that agree with HDF\\,130 are included in Table~\\ref{tbl2}. Even at $151$~MHz the ghost source may not be observable. However, the models predict that the source will be observable in the $74$~MHz band, with luminosity on the order of $10^{43}$~erg~s$^{-1}$, or, equivalently, $10^{20}$~W~Hz$^{-1}$~sr$^{-1}$. New low-frequency observations of HDF\\,130 could test this.  Some models predict the lobe lengths in the lower limit of the error tolerance during the observational congruent window of HDF\\,130, while predicting the other features accurately. If such is the case, perhaps the surrounding density profile is not as simple as we have assumed it to be and allows for the lobes to grow larger than in our models while staying bright in the X-ray. It is plausible that the source is expanding into a pre-existing lobe from a previous episode of jet activity, which cleared away some of the surrounding material and would mean expansion losses are smaller and the lobes can grow larger and brighter.  Importantly, the minimum Lorentz factor of injected particles into the lobe for HDF\\,130 is found to be on the order of $\\gamma_{\\rm min}=1000$ rather than $\\gamma_{\\rm min}=1$. Even a $\\gamma_{\\rm min}=30$ or $\\gamma_{\\rm min}=100$ is not preferred by HDF\\,130: repeating fitting the observable properties of HDF\\,130 with models that have $\\gamma_{\\rm min}=30$ and $\\gamma_{\\rm min}=100$ gives only $4$ congruent models observational windows of at most $3$~Myr, reported in Table~\\ref{tbl1c}.  In the model, a higher $\\gamma_{\\rm min}$ (while keeping injected electron energy density constant) will produce brighter sources without affecting the lobe growth, which is determined by the jet power and the surrounding density profile. Increasing the jet power makes the jet grow larger and brighter. It is the combination of  HDF\\,130 lobe size, which is not exceptionally large, and X-ray brightness which forces the model to require $\\gamma_{\\rm min}\\sim 1000$ to agree with the observational features.    \\begin{table} \\centering  \\caption{Additional models investigated that agree with HDF\\,130 observations}\\label{tbl1c} \\begin{tabular}{@{}lccccccc@{}} \\hline\\hline \\noalign{\\smallskip} \\# & $Q_{\\rm j}^{\\rm a}$ & $t_{\\rm j}^{\\rm b}$ & $\\gamma_{\\rm min}$ & $p$ & $\\beta$ & $\\rho_0^{\\rm c}$ & $t_{\\rm obs}$ ($\\times t_{\\rm j}$) \\\\[2pt] \\hline \\noalign{\\smallskip} [1] & $50$ & $50$ & $30$ & $2.14$ & $1.5$ & $16.7$ & $1.16$ to $1.21$\\\\[2pt] [2] & $10$ & $100$ & $100$ & $2.14$ & $1.5$ & $16.7$ & $1.09$ to $1.12$\\\\[2pt] [3] & $20$ & $50$ & $100$ & $2.14$ & $1.5$ & $16.7$ & $1.16$ to $1.20$\\\\[2pt] [4] & $50$ & $50$ & $100$ & $2.5$ & $1.5$ & $16.7$ & $1.15$ to $1.18$\\\\[2pt] \\hline \\noalign{\\smallskip} \\multicolumn{8}{l}{$^{\\rm a}$ ($\\times10^{38}$ W) \\,\\,\\, $^{\\rm b}$ (Myr) \\,\\,\\, $^{\\rm c}$ ($\\times10^{-23}~{\\rm kg}~{\\rm m}^{-3}$)}  \\\\[2pt] \\end{tabular} \\end{table} The minimum injected Lorentz factor for 6C\\,0905+3955 is also found to be most likely on the order of $\\gamma_{\\rm min}=1000$ (or also marginally $1$), based on only best matching the total lobe luminosities predicted by the model to the observed luminosities. Considering only the models where $p$ is steeper than implied by $\\Gamma$ (these values are bold in Table~\\ref{tbl1b}), we see that only $\\gamma_{\\rm min}=1000$ is preferred. In previous observations, no X-ray emission is seen in a hotspot of 6C\\,0905+3955 (other than highly energetic X-ray synchrotron requiring extremely high Lorentz factors  \\cite{2008MNRAS.386.1774E}) which suggests that there is a low-energy cutoff of the freshly injected particles into the lobe above the $\\gamma\\sim10^3$ particles required for X-ray emission from upscattering on the CMB. Likely, the minimum energy cutoff is just above the critical Lorentz factor which would result in X-ray emission from the hotspot. It is important to note that 6C\\,0905+3955 may be more complicated than described by our simple model, because 6C\\,0905+3955 is asymmetric, probably due to an asymmetric surrounding environment. There may also be complex mechanisms happening in the lobes, such as reflected shocks or interruptions of the jet at the hotspot \\citep{1995MNRAS.277..995L}, which are so far unaccounted for by our model.  The chosen value of $\\gamma_{\\rm min}$ varies by orders of magnitude in previous papers, as it often has to be estimated. The minimum Lorentz factor is assumed to be typically $1$ in previous models of FR~II evolution by \\cite{1997MNRAS.292..723K}, \\cite{1997MNRAS.286..215K}, \\cite{1999AJ....117..677B} and \\cite{2010MNRAS.407.1998N}. \\cite{2005ApJ...626..733C} use a value of $10$, \\cite{1991ApJ...383..554C} use a value of $100$, and \\cite{1998Natur.395..457W} use a value of $1000$. If HDF\\,130 and 6C\\,0905+3955 are typical sources, it may be the case that the minimum energy of particles injected into the lobes is large. The value of $\\gamma_{\\rm min}$ may at first appear as an eclectic, unimportant detail, but \\citet{2011MNRAS.413.1107M} show that the typical value of $\\gamma_{\\rm min}$ can significantly affect estimates for the total population of FR~II sources from a radio luminosity function as it changes the time sources fall below a given flux limit in their evolution. A higher $\\gamma_{\\rm min}$ will also increase the detectability of IC ghosts."
    },
    "1107/1107.3016_arXiv.txt": {
        "abstract": "{NGC 6791 is a  well-studied, metal-rich open cluster that it is so  close to us that  can be imaged  down to luminosities fainter  than that  of the  termination of  its  white-dwarf cooling sequence, thus allowing for an in-depth study of its white dwarf population.} {White dwarfs  carry important information  about the history of  the cluster.   We  use observations  of the  white-dwarf cooling  sequence to constrain  important properties  of the cluster  stellar  population, such  as  the  existence of  a putative population of massive helium-core white dwarfs, and the properties  of a  large population of  unresolved binary white dwarfs.   We also investigate the use  of white dwarfs to disclose  the presence  of cluster subpopulations  with a different  initial chemical  composition, and  we  obtain an upper  bound  to the  fraction  of hydrogen-deficient  white dwarfs.} {We  use  a Monte  Carlo  simulator  that employs  up-to-date evolutionary  cooling   sequences  for  white   dwarfs  with hydrogen-rich   and  hydrogen-deficient   atmospheres,  with carbon-oxygen and  helium cores.  The  cooling sequences for carbon-oxygen  cores account  for the  delays  introduced by both  $^{22}$Ne sedimentation  in  the liquid  phase and  by carbon-oxygen phase separation upon crystallization.} {We  do  not find  evidence  for  a  substantial fraction  of helium-core white dwarfs, and hence our results support the suggestion  that  the origin  of  the  bright  peak of  the white-dwarf luminosity function can only be attributed to a population  of unresolved  binary white  dwarfs.  Moreover, our  results indicate  that if  this hypothesis  is  at the origin  of  the bright  peak,  the  number distribution  of secondary masses  of the population  of unresolved binaries has  to increase  with  increasing mass  ratio between  the secondary and primary  components of the progenitor system. We also find that  the observed cooling sequence appears to be   able   to  constrain   the   presence  of   progenitor subpopulations with different chemical compositions and the fraction of hydrogen-deficient white dwarfs.} {Our simulations  place interesting constraints  on important characteristics of  the stellar populations  of NGC~6791. In particular, we find that  the fraction of single helium-core white dwarfs must be  smaller than 5\\%, that a subpopulation of stars with zero metallicity  must be $\\la 12$\\%, while if the adopted  metallicity of  the subpopulation is  solar the upper limit is  $\\sim 8$\\%.  Finally, we also  find that the fraction   of  hydrogen-deficient   white  dwarfs   in  this particular cluster is surprinsingly small ($\\la 6\\%$).} ",
        "introduction": "\\label{intro}  White dwarfs are the  most common stellar evolutionary endpoint. Their evolution can be described as a very slow cooling process.  Given that these fossil  stars are abundant  and long-lived objects,  they convey important  information  about  the  formation  and  evolution  of  all Galactic populations (Hansen \\&  Liebert 2003; Althaus et al.  2010a). In particular, white dwarfs can  be employed as reliable cosmic clocks to infer  the ages 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) and the system  of globular and open clusters (Kalirai  at al.  2001; Hansen et  al.  2002).  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; Garc\\'\\i a--Berro et al. 2011).  The use of white dwarfs as cosmochronometers relies on the accuracy of the  theoretical cooling  sequences.  Recently,  Garc\\'\\i  a--Berro et al. (2010)  have demonstrated  that the slow  down of the  white dwarf cooling  rate  owing  to  the  release of  gravitational  energy  from $^{22}$Ne sedimentation  (Bravo et al.   1992; Bildsten \\&  Hall 2001; Deloye  \\&  Bildsten 2002)  and  carbon-oxygen  phase separation  upon crystallization  (Garc\\'\\i a--Berro  et al.   1988a; Segretain  et al. 1994) is of fundamental importance to reconcile the age discrepancy of the  very old,  metal-rich  open cluster  NGC~6791.   This raises  the possibility  of using  the  white-dwarf  luminosity  function of  this cluster to constrain its fundamental properties.  NGC~6791  is one  of the  oldest Galactic  open clusters  --  see, for instance,  Bedin  et al.   (2005)  and  Kalirai  et al.   (2007),  and references therein -- and  it is so close to us that  it can be imaged down to  very faint luminosities.   A deep luminosity function  of its well populated white dwarf  sequence has been recently determined from HST observations (Bedin et al.  2008a) and displays a sharp cut-off at low  luminosities, caused by  the finite  age of  the cluster,  plus a secondary  peak at  larger  luminosities, most  likely  produced by  a population of  unresolved binary white  dwarfs (Bedin et  al.  2008b). These characteristics  make this cluster  a primary target to  use the white dwarf cooling sequence to constrain the presence of a population of  massive helium-core  white dwarfs,  the properties  of  the binary white   dwarf  population,  the   hypothetical  presence   of  cluster subpopulations   of  different  metallicity,   and  the   fraction  of hydrogen-deficient  (non-DA)  white  dwarfs.   Here we  address  these issues by  means of  Monte Carlo based  techniques aimed  at producing synthetic  color-magnitude  diagrams   and  luminosity  functions  for NGC~6791 white  dwarfs, which can  be compared with  the observational data.  Firstly, we investigate in detail  the nature of the secondary peak of the white-dwarf luminosity function.  This peak has been atributted to a population of unresolved binary white dwarfs (Bedin et al. 2008b) or to the  existence of a  population of single helium-core  white dwarfs (Hansen  2005).  This  is  a  crucial question  because  if the  first hypothesis is true,  the amplitude of the secondary  peak is such that leads to an unusual fraction  of binary white dwarfs, thus challenging our understanding of the physical processes that rule the formation of binary white dwarfs,  especially at high metallicities.  Consequently, we also study other explanations for the existence of a secondary peak in  the  white-dwarf luminosity  function,  like  the  existence of  a population of  single helium-core white dwarfs.   This explanation was first  put forward by  Hansen (2005)  and Kalirai  et al.   (2007) and later  was challenged,  among  others,  by van  Loon  et al.   (2008). However,  as  we  will  show  below, this  explanation  results  in  a white-dwarf luminosity function that is at odds with the observed one, and hence the most likely explanation for the secondary bright peak is the  population of  unresolved binaries.   Indeed, there  is  a strong reason to suspect  that the bright peak of  the white-dwarf luminosity function of  NGC~6791 is caused  by a population of  unresolved binary white dwarfs, namely, that the two peaks of the white-dwarf luminosity function are separated by $0.75^{\\rm  mag}$. This is just exactly what it  should be  expected if  the bright  peak is  caused by  equal mass binaries.   Hence, if  this explanation  for  the bright  peak of  the white-dwarf luminosity function  is true this, in turn,  enables us to study the  properties of the  population of such binary  white dwarfs. Specifically, we study how different distributions of secondary masses of  the  unresolved binary  white  dwarfs  affect the  color-magnitude diagram and the white-dwarf luminosity function.  Additionally, we test whether  the white-dwarf luminosity function can provide   an  independent   way  to   check  for   the   existence  of subpopulations  within  a  stellar  system.   The  presence  of  these subpopulations has been found in several Galactic globular clusters of which $\\omega$ Cen is,  perhaps, the most representative one (Calamida et al.   2009).  The  appearance of NGC~6791  color-magnitude diagram, and the lack of any  significant chemical abundance spread (Carraro et al.     2006)    points   toward    a    very   homogeneous    stellar population. However,  a recent paper  by Twarog et al.   (2011) raises the possibility  that there  could be a  1~Gyr age  difference between inner  and outer  regions of  the cluster.   Nevertheless,  within the field covered by  the white dwarf photometry used  in our analysis the age difference should  be negigible.  This is different  from the case of  individual  Galactic globular  clusters  that  host  at least  two distinct  populations with  approximately the  same age  and different abundance  patterns of  the C-N-O-Na  elements  -- see  Milone et  al. (2010) for   a  recent   brief  review.    Taking  advantage   of  the well-populated  cooling  sequence   in  the  observed  color-magnitude diagram,   we  test   whether   modeling  the   cluster  white   dwarf color-magnitude  alone  can  exclude  the presence  of  subpopulations generated  by progenitors with  a metallicity  different from  the one measured spectroscopically.  As a final study, we use the luminosity function of NGC~6791 (Bedin et al.   2008a),  and  the  fact  that white  dwarfs  with  hydrogen-rich atmospheres  (of  the DA  type)  and non-DA  white  dwarfs  cool at  a different rate, to place constraints on the fraction of these objects. This is  an important point because  Kalirai et al.  (2005) have found that in  the young  open star  cluster NGC~2099 there  is a  dearth of non-DA  white dwarfs.   These authors  attributed the  lack  of non-DA objects to  the fact  that possibly the  hot, high-mass  cluster white dwarfs  -- white  dwarfs  in this  cluster  are estimated  to be  more massive than  field objects  in the same  temperature range --  do not develop sufficiently extended helium  convection zones to allow helium to be brought to the surface and turn a hydrogen-rich white dwarf into a helium-rich one.  Studying the  fraction of non-DA white dwarfs in a different  open cluster  could  provide additional  insight into  this question.  Moreover, Kalirai et al.  (2007) analyzed a sample of $\\sim 15$ white  dwarfs in NGC  6791, and although  the sample was  far from being complete, all them were of the DA type.  The  paper  is organized  as  follows.   Sect.~2  summarizes the  main ingredients of our  Monte Carlo code plus the  other basic assumptions and procedures necessary to  evaluate the characteristics of the white dwarf   population    for   the   different    simulations   presented here. Specifically, we discuss  the most important ingredients used to construct  reliable  color-magnitude  diagrams and  the  corresponding white-dwarf luminosity functions. Sect.~3 presents the main results of our  simulations for  all points  already mentioned  in  this section. Finally, Sect.~4 closes the paper with a summary of our conclusions. ",
        "conclusions": "\\label{concl}  In this paper we have investigated several important properties of the stellar  population  hosted  by  the  very  old  (8  Gyr),  metal-rich ([Fe/H]$\\simeq  0.4$) open  cluster NGC~6791.   This cluster  has been imaged  below the  luminosity of  the termination  of its  white dwarf cooling  sequence   (Bedin  et  al.   2005,   2008a).   The  resulting white-dwarf luminosity  function enables us not only  to determine the cluster  age (Garc\\'\\i  a--Berro et  al.  2010),  but  other important properties as  well.  Among  these, we mention  the properties  of the population of unresolved binary white dwarfs, the existence of cluster subpopulations, and the fraction of non-DA white dwarfs.  The origin of  the bright peak of the  white-dwarf luminosity function was  investigated,   exploring  in  detail   the  alternative  massive helium-core white  dwarf scenario.  Our  conclusion is that  this peak cannot  be attributed  to  a population  of  single helium-core  white dwarfs.  The  more realistic possibility left to  explain this feature is  a  population  of  unresolved  binary  white  dwarfs.   This  huge population of unresolved  binary white dwarfs has allowed  us to study the properties of the parent  population.  To this purpose, we studied the properties of  the distribution of secondary masses  in the binary progenitor  system  and  its  effects in  the  white-dwarf  luminosity function.   Specifically, we  tested four  different  distributions of secondary masses and  we found that only those  distributions that are monotonically increasing  with the mass ratio are  consistent with the observational  data.  Additionally, as  a test  case, we  verified the ability of the white-dwarf luminosity function to assess the existence of subpopulations  within a  stellar system.  We  have found  that the presence of a $Z=0$ subpopulation is inconsistent with the white-dwarf luminosity function,  the maximum fraction  allowed by the  data being 12\\%.  If the metallicity of the subpopulation is solar, this fraction is 8\\%.  Finally, we  found that  the fraction of  non-DA white dwarfs  in this cluster is  unusually small,  on the  order of 6\\%  at most,  and much smaller than  the corresponding one  for field white dwarfs,  which is $\\sim 20\\%$. This shortage of  non-DA white dwarfs is a characteristic shared with  another open cluster, NGC~2099.  However,  the deficit of non-DA white dwarfs is even higher in the case of NGC~2099, given that for this  cluster recent exhaustive observations have  found no single white dwarf of the non-DA type (Kalirai et al.  2005)."
    },
    "1107/1107.4894_arXiv.txt": {
        "abstract": "{Products from massive-star nucleosynthesis have been measured with SPI on INTEGRAL: Characteristic gamma-ray lines from radioactive decays of long-lived $^{26}$Al and $^{60}$Fe isotopes, and from $^{44}$Ti decay ($\\tau$=89y). Detections of both these isotopes has laid the foundation to peek into massive-star interiors, through different views at those measurements. The gamma-ray flux can be converted into amounts of these radioactive isotopes, but the constraints which derive from this for massive star models involve additional steps. -- Earlier results had demonstrated the basic constraints inherent to such radioactivity data, i.e. detection of $^{26}$Al is a calibration for massive-star yields, $^{60}$Fe enables an isotopic-yield ratio test eliminating modeling and observing bias aspects, and $^{44}$Ti searches showed that its production does not occur homogeneously over core-collapse events. The current status of $^{26}$Al and $^{60}$Fe observations and their analysis is reaching a threshold for astrophysical insights: Specific regions of massive star groups and their radioactivity gamma-rays have recently been investigated, such as Sco-Cen and Orion. As for $^{44}$Ti, spectroscopical information constrains ejection velocities of these inner supernova ejecta for the Cas A event. These parameters for massive stars can be related to  other fields of astronomy, of nuclear physics experiments and theory, and of theoretical astrophysics. Thus,  model improvements appear on the horizon, as e.g.  implemented in the EuroGenesis program framework of the European Science Foundation.}  \\FullConference{8$^{th}$ INTEGRAL Science Workshop\\\\ {\\it The Restless Gamma-Ray Universe} \\\\ October 27-30, 2010\\\\ Dublin, Ireland}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "%  The deepening of exposure resulting from the extended mission now gradually lifts details beyond the minimum thresholds of scientific significance, which could be important for improving insights into massive star evolution and their nucleosynthesis. SPI's capabilities to measure the gamma-ray lines associated with the decays of \\Al, \\Fe, and positron annihilations now combines with the modest spatial resolution to allow localized (i.e. region or star-group-specific) studies. These are more precise than global galactic averages, thus helping to constrain our uncertain models of massive star structure and evolution. Specifically, regions along the inner Galactic ridge, such as Scorpius Centaurus, Aquila, but also the Orion region, are promising laboratories for massive star studies with INTEGRAL in its later mission years.  The European Science Foundation's ``EuroGenesis'' program (2010-2013) assembles several groups of scientists working on a varity of stellar-model, nuclear-reaction, and astronomical data collection aspects. This program also includes a special investigation of dust production from massive stars, adding a potential exploitation of infrared and meteoritic data for learning about massive star structure and evolution. It is one example which demonstrates how INTEGRAL measurements have now found their way into studies of general astrophysical interest and broad application.  {\\bf Acknowledgements} The INTEGRAL project of ESA and the SPI project has been completed under the responsibility and leadership of CNES/France.We are grateful to ASI, CEA, CNES, DLR, ESA, INTA, NASA and OSTC for support. The SPI anticoincidence system is supported by the German government through DLR grant 50.0G.9503.0.. The work for this paper was supported by the Munich Cluster of Excellence on ``Origins and Evolution of the Universe''."
    },
    "1107/1107.5360_arXiv.txt": {
        "abstract": "We analyzed the distribution of the RC stars throughout Galactic bulge using 2MASS data. We mapped the position of the red clump in 1~sq.~deg.  size fields within the area $|l|\\le8.5^{\\circ}$ and $3.5^{\\circ}\\le |b| \\le8.5^{\\circ}$, for a total of 170~sq.~deg. The red clump seen single in the central area splits into two components at high Galactic longitudes in both hemispheres, produced by two structures at different distances along the same line of sight. The X-shape is clearly visible in the $Z$--$X$ plane for longitudes close to $l=0^{\\circ}$ axis. Crude measurements of the space densities of RC stars in the bright and faint RC populations are consistent with the adopted RC distances, providing further supporting evidence that the X-structure is real, and that there is approximate front-back symmetry in our bulge fields. We conclude that the Milky Way bulge has an X-shaped structure within $|l|\\lesssim 2^{\\circ}$, seen almost edge on with respect to the line of sight. Additional deep NIR photometry extending into the innermost bulge regions combined with spectroscopic data is needed in order to discriminate among the different possibilities that can cause the observed X-shaped structure. ",
        "introduction": "In the last decade consensus has been reached that the Galactic bulge has a boxy shape in the COBE/DIRBE near infrared map \\citep{1995ApJ...445..716D} because it contains a bar. This bar is tilted by $\\sim 20-25$ degrees from the Sun-Galactic center direction, with axis ratios 1\\,:\\,0.33\\,:\\,0.23, and a scale length of $\\sim 1.5$~kpc \\cite[e.g.,][and references therein]{1994ApJ...429L..73S, 1997ApJ...477..163S, 2005MNRAS.358.1309B, 2007MNRAS.378.1064R}. Most of these studies make use of horizontal branch red clump (RC) stars as tracers of the bulge 3D structure, due to their narrow range in color, brightness, and well calibrated absolute magnitude.  However, recent analysis of the RC stars across the bulge area suggested a more complex X-shaped structure \\citep{2010IAUS..265..279M, 2010IAUS..265..271Z, 2010ApJ...721L..28N, 2010ApJ...724.1491M}. Along to $l=0^{\\circ}$ axis, both above and below the plane, the RC appears to split in two components. Both components become closer and finally merge at the latitude of Baade's Window ($b=-4^{\\circ}$) and below, while they open up and diverge at higher latitudes both above and below the plane. The analysis by \\cite{2011ApJ...730..118N} is based on the RC distribution from the OGLE-II photometry, while \\cite{2010ApJ...724.1491M} analyzed the color magnitude diagrams (CMDs) from 2MASS on a roughly symmetric field grid across the bulge. For the field at $(l,b)=(0^{\\circ},-6^{\\circ})$, in particular, \\cite{2010ApJ...724.1491M} had deeper SOFI@NTT infrared photometry ($K_{\\rm s}\\sim18$~mag, in comparison to $K_{\\rm s}\\sim14.3$~mag from 2MASS) as well as deep optical data from WFI@ESO/MPG telescope ($V\\sim23.0$~mag, $I\\sim22.5$~mag).   Here we complete the analysis of the RC distribution across a much wider bulge area using 2MASS data. The RC was mapped at positive and negative latitudes in 170 fields of 1~sq.~deg. size, within $|l|\\le8.5^{\\circ}$ and $3.5^{\\circ}\\le |b| \\le8.5^{\\circ}$.  Our results are in agreement with previous studies, and extend their findings, showing the presence of an X-shaped structure across the whole Galactic bulge. ",
        "conclusions": "The position of the RC(s) was mapped across a large bulge area, using 2MASS photometry.  By using the RC(s) magnitude as distance indicator, density maps were constructed along the line of sight, for the region $|l|\\le8.5^{\\circ}$ and $3.5^{\\circ}\\le |b|\\le8.5^{\\circ}$.  The main conclusion of the present paper is that in the Milky Way bulge has an X-shape, seen almost edge on, with an inclination of $\\sim 20^\\circ$ with respect to the line of sight. The X is clearly seen within $|l|\\lesssim2^{\\circ}$, with the closer arm remaining at positive latitudes while at negative latitudes only the farther arm is present. The {\\it arms} of the X merge together both above and below the plane, for latitudes $|b|\\le 4^{\\circ}$. Therefore, in Baade's Window at $b=-4^{\\circ}$ one would just see an elongated structure, corresponding to the two lower arms being {\\it almost} completely merged.  This feature has been classically interpreted as the Galactic bar.  Our crude measurement of the space densities of RC stars as a function of vertical height below the Galactic plane is consistent with the adopted RC star distances, if the bulge structures studied here are roughly symmetric from front to back.  This removes any lingering doubt about the interpretation of the RC $K$-band magnitudes as due to distance, and strengthens the conclusion that the bulge contains an X-shape structure.  Our densities are slightly higher for the background structure; this possibility should be investigated with more sophisticated techniques.  Several galaxies are known to host X-shaped bulges (e.g., NGC~128, NGC~3625, NGC~4469), and some numerical simulations predict their formation, with a particularly strong buckling of a bar \\citep{2002MNRAS.337..578P, 2005MNRAS.358.1477A, 1995ApJ...447L..87M}. \\cite{2010ApJ...724.1491M} review some of the proposed mechanisms to form such structures.  The X-shape structure seems to be relatively concentrated, with respect to other galaxies (e.g., NGC~4710) and probably is the dominant component in the inner 2~kpc.  It remains to be understood how the bulge 3D structure can be reconciled with the other features, such as the presence of a radial metallicity gradient \\citep{2008A&A...486..177Z, 2010ApJ...725L..19B}, the kinematic difference between the metal poor and metal rich component in Baade's Window \\citep{2010A&A...519A..77B} and the behavior of the alpha element ratios, which seems to be constant in different fields (Gonzalez et al. 2011). It should be kept in mind, however, that all the properties listed above have been derived from observations on $l\\simeq0^{\\circ}$ only. Just as the shape turned out to be much more complex as soon as we analyzed a larger area, the other properties might turn out to be a partial view of the global picture.  Radial velocities of stars in the two RCs, at $b=-8^{\\circ}$ revealed no differences \\citep{2011arXiv1104.0223D}. However a similar analysis at $b=-6^{\\circ}$ reveals the presence of two peaks at $v_r=+100~/-100~$km\\,s$^{-1}$ in the faint/bright RC, respectively (Vasquez et al.  2011, in preparation). Proper motions also do not seem to differ between stars in the two clumps \\citep{2007AJ....134.1432}.  Further investigations are clearly necessary to completely understand our bulge structure, kinematics, chemistry and, ultimately, origin. The present mapping traces the X-shaped structure across the whole bulge area and may help to guide further analysis, as well as to design future observations of the Milky Way bulge."
    },
    "1107/1107.5199.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We seek to constrain the formation of the Galactic bulge by means of analysing the detailed chemical composition of a large sample of red clump stars in Baade's window. } % methods heading (mandatory) {We measure [Fe/H] in a sample of 219 bulge red clump stars from R=20000 resolution spectra obtained with FLAMES/GIRAFFE at the VLT, using an automatic procedure, differentially to the metal-rich local reference star \\MuLeo. For a subsample of 162 stars, we also derive [Mg/H] from spectral synthesis around the \\ion{Mg}{I} triplet at $\\rm\\lambda$ 6319 \\AA. } % results heading (mandatory) {The Fe and Mg metallicity distributions are both asymmetric, with median values of $+0.16$ and $+0.21$ respectively. The iron distribution is clearly bimodal, as revealed both by a deconvolution (from observational errors) and a Gaussian decomposition. The decomposition of the observed Fe and Mg metallicity distributions into Gaussian components yields two populations of equal sizes (50\\% each): a metal-poor component centred around [Fe/H]~$\\rm=-0.30$ and [Mg/H]~$\\rm=-0.06$ with a large dispersion and a narrow metal-rich component centred around $\\rm [Fe/H]=+0.32$ and $\\rm [Mg/H]=+0.35$. The metal poor component shows high [Mg/Fe] ratios (around 0.3) whereas stars in the metal rich component are found to have near solar ratios. Babusiaux et al. (2010) also find kinematical differences between the two components: the metal poor component shows kinematics compatible with an old spheroid whereas the metal rich component is consistent with a population supporting a bar. In view of their chemical and kinematical properties, we suggest different formation scenarii for the two populations: a rapid formation timescale as an old spheroid for the metal poor component (old bulge) and for the metal rich component, a formation over a longer time scale driven by the evolution of the bar ({\\em pseudo-bulge}). } % conclusions heading (optional), leave it empty if necessary { Guided by these results, we build a simple model combining two components: a simple closed box model to predict the metal poor population contribution, whereas the metal rich population is modelled using the observed local thin disc metallicity distribution, shifted in metallicity. The {{\\em pseudo-bulge}} is compatible with being formed from the inner thin disc, assuming large (but plausible) values of the gradients in the early Galactic disc. } ",
        "introduction": "The Milky-Way bulge has been the subject of quite intense debates in the community, as its status is not yet fully established. With various stellar population characteristics similar to those of the central old spheroids found at the centre of earlier-type galaxies, and others (mostly geometrical and kinematical) that rather point towards a very strong influence of a bar responsible for a {\\em pseudo-bulge}, it sits at the border between these two types of bulges. The formation scenarios for bulges can be classified in three different types: (i) initial collapse of gas at early times \\citep[see e.g. ][]{Eggen1962}; (ii) merging subclumps, either through an early disc evolution \\citep{Noguchi1999,Immeli2004}, or through mergers \\citep{Aguerri2001,Scannapieco2003,Nakasato2003}; (iii) secular evolution of the Galactic disc \\citep{Combes1981,Pfenniger1990,Raha1991}. The merger scenario itself is similar to an early collapse scenario from the point of view of the formation characteristic timescale (early and fast). It was also recently suggested that the Galactic bulge could be the result of both formation processes, with an {\\it old spheroid} complemented by a {\\em bar-driven pseudo-bulge} \\citep{Nakasato2003,Kormendy2004,Gerhard2006}. The relative importance of the two processes (or even populations) however remains to be established.  The presence of a bar in the inner Galaxy has been suggested by \\citet{deVaucouleurs1964} from gas kinematics and confirmed since then by numerous studies including infrared surface brightness map \\citep[e.g.][]{2MASS_survey}, star-counts, red-clump distances \\citep[e.g.][]{Babusiaux2005,Nishiyama2006,Rattenbury2007a}, microlensing and stellar kinematics \\citep[e.g.][]{Zhao1994,Howard2008,Howard2009}. The boxy aspect of the bulge, detected in the infrared light profile \\citep[e.g.][]{Dwek1995}, also argues for a {\\em pseudo-bulge} secularly evolved from the galactic disc.   %% Wide near infrared surveys such as COBE/DIRBE \\citep{Dwek1995}, DENIS %% \\citep{DENIS_survey} and 2MASS \\citep{2MASS_survey} clearly revealed %% the boxy-shaped nature of the Galactic bulge, highlighting its %% bar-like nature. Several authors confirmed and refined the angle and %% length of the bar, based on star-counts \\citep[e.g.][]{Babusiaux2005}. %% Kinematics of the inner galaxy also seem to point towards barlike %% motions \\citep[e.g.][]{Zhao1994,Howard2008,Howard2009}, clearly pointing %% towards a {\\em pseudo-bulge} secularly evolved from the disc.  On the other hand, photometric studies in selected windows on the Galactic bulge, in the visible and the near infrared, soon made apparent that the stellar populations in the bulge is old and metal-rich \\citep{Ortolani1995,Feltzing2000,Kuijken2002,Zoccali2003,Clarkson2008}. Spectroscopically, \\citet{Rich1988} obtained one of the first metallicity distribution in Baade's window based on low-resolution spectra of M giants, followed by \\citet{Ibata1995a,Ibata1995b} using K giants and \\citet{Sadler1996} using red clump stars, all finding a large metallicity dispersion. High resolution spectra of a limited number of stars (10 to 20) in Baade's window \\citep{McWilliam1994,Fulbright2006,Fulbright2007}, and more recently by our group \\citep{Zoccali2006,Lecureur2007} in a larger sample of $~$50 stars in four windows of the Galactic bulge, showed enhanced [$\\alpha$/Fe], compatible with a fast chemical enrichment of the Galactic bulge. Both the chemical and age properties of stellar popultions in the Galactic bulge thus point towards a rapid bulge formation. Early combined metallicity and kinematics measurements also pointed towards bulge formation through dissipational collapse \\citep{Minniti1996}.  More recently still, our group presented in \\citet{Zoccali2008} the first metallicity distribution entirely based on high-resolution spectra of large sample of red giant branch stars in three fields of the Galactic bulge (close to the minor axis, at b$\\sim -4, -6$ and $-12^{\\circ}$). In this paper we will show how determining metallicity distributions from a large ($> 200$) and almost uncontaminated sample of red clump stars in Baade's window, observed at high spectral resolution can lead to significant improvements in our understanding of the origin and nature of the Galactic bulge. In particular, we show how using two independent elements (iron {\\em and} magnesium) as metallicity tracers (and their ratio [Mg/Fe]) reveals the nature of the stellar population. In Sect.~2, we present our target selection and observations with FLAMES on VLT, in Sect.~3 we detail the stellar parameter and elemental abundance measurement methods, while Sect.~4 examines the issues of sample representativity and contamination. Finally, in Sect.~5, we discuss the resulting [Fe/H] and [Mg/H] metallicity distributions and [Mg/Fe] trends, show that the sample can be separated into two distinct populations and interpret these in the framework of various formation scenarios and chemical evolution models. Sect.~6 gives our conclusions from this work. ",
        "conclusions": "We have analysed a sample of 219 bulge red clump stars in Baade's window from R=20000 resolution spectra obtained with FLAMES/GIRAFFE at the VLT. This sample has been selected based on the location of stars in the Colour-Magnitude Diagram, in order to minimise the contamination by other Galactic components ($\\rm<10$ \\% from the latest population synthesis Besan\\c{c}on model \\citep{Robin2003,Picaud2004}, and hence is a good tracer of the bulge metallicity distribution (MD) in the Baade's window.  For these stars, with an automatic procedure, we have derived the stellar parameters and Fe abundances using an iron linelist established differentially to \\MuLeo\\ and for a subsample of 162 stars (built by excluding stars with high uncertainties on [Fe/H]), we also derived the Mg abundances with spectral synthesis around the \\ion{Mg}{I} triplet at $\\rm\\lambda$ 6319 \\AA.  The Fe and Mg metallicity distributions are both asymmetric, with mean values of $\\rm+0.05 \\pm 0.03$ and $\\rm+0.14 \\pm 0.03$, and median values of $\\rm0.16$ and $\\rm0.21$ respectively. They show a small proportion of stars at low metallicities ($\\rm<3$ \\% with $\\rm[Fe/H]<-0.7$ or $\\rm[Mg/H]<-0.4$), only extending down to $\\rm[Fe/H]=-1.1$ or $\\rm[Mg/H]=-0.7$. The iron MD is not well reproduced by the predictions of the chemical models of the bulge available in the literature \\citep{Ballero2007, Ballero2007BIS}, neither by models adding dynamical aspects to chemical prescriptions such as the one of \\citep{Immeli2004}.  The decomposition of the observed Fe and Mg MDs into Gaussian components has revealed the presence of two populations of equal sizes (50\\% each): a metal-poor component centred around [Fe/H]~$\\rm=-0.30$ and [Mg/H]~$\\rm=-0.06$ with a large dispersion and a narrow metal-rich component centred around $\\rm [Fe/H]=0.32$ and $\\rm [Mg/H]=+0.35$. The metal poor component shows high [Mg/Fe] ratios (around 0.3) whereas stars in the metal rich component are found to have near solar ratios. \\citet{Babusiaux2010} also found kinematical differences between the two components: the metal poor component shows kinematics compatible with an old spheroid whereas the metal rich component is consistent with a population supporting a bar. In view of their chemical and kinematical properties, we suggest different formation scenarios for the two populations: a rapid formation timescale as an old spheroid for the metal poor component (old bulge) and for the metal rich component, a formation over a longer time scale driven by the evolution of the bar ({\\em pseudo-bulge}). This latter component represents ~50\\% of the total Baade's Window population, and is of old age \\citep[ as indicated by deep colour-magnitude diagrams, see e.g.]{Clarkson2008}, indicating that a fairly massive disc was present some 10Gyr ago, which may turn out to be a strong constraint on models of disc formation. It is suggestive that the local thick disc and old bulge display the same [Mg/Fe] enhancement, pointing at similarly rapid star formation process. At this point, it is however not possible to constrain further the relationship that the old bulge bears with the thick disc, both because of lack of constraints on the central parts of the thick disc (both chemistry and kinematics), and because our bulge sample being located on the minor axis, we have a very poor constraint on rotation and the disky nature of this component can therefore not be probed from our data. We however not that in their recent study, \\citet{Bensby2010tkd} finds that the inner parts of the thick disc is chemically very similar to that in the solar neighbourhood, lending strong support to the idea that the metal-poor bulge is indeed chemically related to the thick disc. We also note that the numerous metal-rich globular clusters of the Milky Way, that are concentrated towards the inner parts of the galaxy, would be classified as old bulge in our scheme (see e.g. Bica et al. 2006 that show clusters with $\\rm [Fe/H]<-0.75$ are confined within 2kpc of the centre). It is interesting to note that the most metal-rich globular cluster known so far is NGC 6528 (Ortolani et al. 2007), with $\\rm [Fe/H]=-0.2$ (Zoccali et al. 2004), therefore within the metal-poor bulge component identified here.  Guided by these results, we built a simple model combining two components. We used a simple closed box model to predict the metal poor population contribution, whereas we represented the metal rich population by the MD of the local thin disc \\citep{Fuhrmann2008}, shifted to the centre of the corresponding metal rich population. The resulting combined MD shows a good agreement with the red clump one on the whole metallicity range except for the very metal poor and very metal rich regimes, where more stars are predicted than observed. The bulge (old bulge) thus seems to exhibit the G dwarfs problem. %The sharp cutoff observed a very high [Fe/H] and [Mg/H] values suggest %that the local disc and the inner disc could have had different %stellar formation histories. The formation and evolution of the bar %has most probably influenced the stellar formation in this region. The  shift we had to apply to the local thin disc MD to fit the metal-rich population is however quite large (although not impossible) when compared with plausible values of the gradients in the early Galactic disc. One possibility could be that these stars were born in a pre-enriched environment, enriched by the remnant of the old bulge.  The scenario(s) proposed here should now be turned into proper modelling (chemical and kinematical) of the bulge that would include two formations events (linked or not), to evaluate their viability. Expanding similar combined chemical and kinematical studies of large samples towards various directions in the Galactic bulge would also be of great help for answering these questions. In particular, probing the bulge at lower latitudes and along longitude should trace the relative importance of the metal-rich and metal-poor populations and definitively establish their proposed different (kinematical and chemical) nature as an old bulge and {\\em pseudo-bulge}. Because these regions are much more heavily obscured than Baade's window, this kind of observations will have to be carried out in the infrared. Finally, we note that our group has undertaken similar red clump observations of the inner parts of the Galactic disc that are expected to shed new light on the true relation between the disc and the {\\em pseudo-bulge. }"
    },
    "1107/1107.5226_arXiv.txt": {
        "abstract": "The nonlinear dynamics of electron-acoustic localized structures in a collisionless and unmagnetized plasma consisting of ``cool'' inertial electrons, ``hot'' electrons having a kappa distribution, and stationary ions is studied. The inertialess hot electron distribution thus has a long-tailed suprathermal (non-Maxwellian) form. A dispersion relation is derived for linear electron-acoustic waves. They show a strong dependence of the charge screening mechanism on excess suprathermality (through $\\kappa$). A nonlinear pseudopotential technique is employed to investigate the occurrence of stationary-profile solitary waves, focusing on how their characteristics depend on the spectral index $\\kappa$, and the hot-to-cool electron temperature and density ratios. Only negative polarity solitary waves are found to exist, in a parameter region which becomes narrower as deviation from the Maxwellian (suprathermality) increases, while the soliton amplitude at fixed soliton speed increases. However, for a constant value of the true Mach number, the amplitude decreases for decreasing $\\kappa$. ",
        "introduction": "}  Electron-acoustic waves may occur\\ in plasmas characterized by a co-existence of two distinct electron populations, here referred to as \\textquotedblleft cool\\textquotedblright\\ and \\textquotedblleft hot\\textquotedblright% \\ electrons. These are electrostatic waves of high frequency (in comparison with the ion plasma frequency), propagating at a phase speed which lies between the hot and cool electron thermal velocities. On such a fast (high frequency) dynamical scale, the positive ions may safely be assumed to form a uniform stationary charge background simply providing charge neutrality, yet playing no essential role in the dynamics. The cool electrons provide the inertia necessary to maintain the electrostatic oscillations, while the restoring force comes from the hot electron pressure.  A matter of importance in electrostatic wave propagation (although inevitably overlooked in fluid plasma models) is Landau damping, which becomes stronger when the phase velocity approaches the thermal velocity of either electron component, thus the wave can propagate in the plasma only within a restricted range of parameter values. It turns out that electron-acoustic waves are weakly damped for a temperature ratio $T_{c}/T_{h} \\lesssim0.1$ and provided that the cool electrons represent an intermediate fraction of the total electron density: $0.2\\lesssim n_{c}/(n_{c}+n_{h})\\lesssim0.8$% .~\\cite{Tokar1984,Gary1985,Mace1990,Berthomier1999} The wavenumber $k$ to minimize damping lies roughly between $0.2\\lambda_{Dc}^{-1}$ and $0.6\\lambda_{Dc}^{-1}$ (where $\\lambda_{Dc}$ is the cool electron Debye length). These results on bi-Maxwellian plasmas were later extended to include the effect of the excess suprathermality of the hot electrons \\cite{Mace1999} (a physical feature to be discussed below). It was found that excess suprathermal electrons do cause a modification of the damping curves, but the overall qualitative conclusion remains unchanged: electron-acoustic waves survive Landau damping over a wide range of parameter values.~\\cite{Mace1999} However, care must be taken in the choice of plasma configuration when studying nonlinear electron-acoustic structures, so as to ensure that one is considering a region of parameter space in which Landau damping is minimized.  Electron-acoustic waves occur in laboratory experiments \\cite{Henry1972,Ikezawa1981} and space plasmas, e.g., in the Earth's bow shock \\cite{Thomsen1983,Feldman1983,Bale1998} and in the auroral magnetosphere \\cite{Tokar1984,Lin1984}. They are associated with Broadband Electrostatic Noise (BEN), a common high-frequency background activity, regularly observed by satellite missions in the plasma sheet boundary layer (PSBL) \\cite{Matsumoto1994,Cattell1999,Kakad2009}. BEN emission includes a series of isolated bipolar pulses, within a frequency range from $\\sim10$ Hz up to the local electron plasma frequency ($\\sim10$ kHz) \\cite{Matsumoto1994}. This clearly suggests that BEN is related to electron dynamics rather than to the ions \\cite{Matsumoto1994,Kakad2009}.  In the standard bi-Maxwellian picture, the two electron species would each be assumed to be in a (different) thermal Maxwellian distribution, parameterized via two distinct temperature values, $T_{c}$ and $T_{h}$, respectively \\cite{Watanabe1977, Nishihara1981, Berthomier1998}. Contrary to this picture, space and laboratory plasmas often possess an excess population of suprathermal electrons, a fact which is reflected in a power law distribution at high velocity (above the electron thermal speed). This excess suprathermality phenomenon is well modeled by a generalized Lorentzian or $\\kappa$-distribution \\cite{Summers1991,Hellberg2000,Baluku2008,Hellberg2009,Pierrard2010}. The common form of the isotropic (three-dimensional) generalized Lorentzian or $\\kappa$-distribution function is given by \\cite{Summers1991,Baluku2008,Hellberg2009}% \\begin{equation} f_{\\kappa}(v)=n_{0}(\\pi\\kappa\\theta^{2})^{-3/2}\\frac{\\Gamma(\\kappa+1)}% {\\Gamma(\\kappa-\\frac{1}{2})}\\left(  1+\\frac{v^{2}}{\\kappa\\theta^{2}}\\right) ^{-\\kappa-1}, \\label{eq1_60}% \\end{equation} where $n_{0}$ is the equilibrium number density of the electrons, $v$ the velocity variable, and $\\theta$ the most probable speed, which acts as a characteristic ``modified thermal speed\", and is related to the usual thermal speed $v_{th,e}=(2k_{B}T_{e}/m_{e})^{1/2}$ by $\\theta=v_{th,e}\\left[  (\\kappa -\\tfrac{3}{2})/\\kappa\\right]  ^{1/2}$. Here $k_{B}$ is the Boltzmann constant, $m_{e}$ the electron mass and $T_{e}$ the temperature of an equivalent Maxwellian having the same energy content. \\cite{Hellberg2009} The term involving the Gamma function ($\\Gamma$) arises from the normalization of $f_{\\kappa}(v)$, viz., $\\int f_{\\kappa}(v)d^{3}v=n_{0}$. Here, suprathermality is denoted by the spectral index $\\kappa$, with $\\kappa> \\frac{3}{2}$, for reality of the most probable speed, $\\theta$.~\\cite{Hellberg2009} Low values of $\\kappa$ are associated with a significant number of suprathermal particles; on the other hand, for $\\kappa\\rightarrow\\infty$ a Maxwellian distribution is recovered.  The $\\kappa$-distribution was first applied to model velocity distributions observed in space plasmas that were Maxwellian-like at lower velocities, but had a power-law form at higher speeds \\cite{Vasyliunas1968}, and was later applied in a variety of studies, successfully fitting many real space observations, e.g., \\cite{Feldman1983,Christon1988,Pierrard1996,Pierrard2010}. Typical $\\kappa$ values usually lie in the range $2<\\kappa<6$. For example, observations in the earth's foreshock satisfy $3<\\kappa_{e}<6$, \\cite{Feldman1983} measurements of plasma sheet electron and ion distributions yield $\\kappa_{i}=4.7$ and $\\kappa_{e}=$ $5.5$ (here, $e$ denotes electrons and $i$ ions), \\cite{Christon1988} and coronal electrons in the solar wind are modeled with $2<\\kappa_{e}<6$ \\cite{Pierrard1996}. Recent observations of the radial distribution of the electron population in Saturn's magnetosphere also point towards a kappa distribution ($\\kappa_{e}\\simeq2.9$ - $4.2)$% ~\\cite{Schippers2008}. Therefore, we focus our interest in the following on the range $2<\\kappa<6$; in fact, the Maxwellian limit is already practically attained for values above $\\kappa\\simeq10$.  A linear analysis of electron-acoustic waves was first carried out by assuming an unmagnetized Maxwellian homogeneous plasma, which exhibited a heavily damped acoustic-like solution in addition to Langmuir waves and ion-acoustic waves. \\cite{Fried1961} Those early results were later extended to take into account the effect of excess suprathermal particles \\cite{Mace1999,Hellberg2002}, whose presence in fact results in an increase in the Landau damping at small wavenumbers, in particular when the hot electron component is dominant \\cite{Mace1999,Baluku2011}. Studies of linear and nonlinear electron-acoustic waves in plasmas with nonthermal electrons have received a great deal of interest in recent years \\cite{Dubouloz1991,Singh2004,Kourakis2004,Cattaert2005,Verheest2005}. Negative potential solitary structures were shown to exist in a two-electron plasma, either for Maxwellian \\cite{Dubouloz1991} or for nonthermal \\cite{Singh2004,Sultana2010} hot electrons. Interestingly, either incorporation of finite inertia~\\cite{Cattaert2005,Verheest2005} or the addition of a beam component~\\cite{Berthomier2000,Mace2001} may lead to the existence of positive and negative potential solitons. A recent investigation has established the properties of modulated electron-acoustic wavepackets in kappa-distributed plasmas, and has studied the effect of suprathermality on the amplitude (modulational) stability.~\\cite{Sultana2011}  In this paper, we study the linear and nonlinear dynamics of electron-acoustic waves in a plasma consisting of cool adiabatic electron and hot $\\kappa $-distributed electrons, in addition to stationary ions. The paper is organised as follows. An electron-plasma-fluid model is presented in Section\\ \\ref{model}. In Section \\ref{DR}, a linear dispersion relation is derived and discussed. In Section \\ref{nonlinear}, the Sagdeev pseudopotential method is employed to investigate the occurrence of stationary profile electrostatic solitary waves. In Section \\ref{existence}, we depict the existence domain of the electron-acoustic solitary waves. Section \\ref{investigation} is devoted to a parametric investigation of the form of the Sagdeev pseudopotential and of the characteristics of electron-acoustic solitary waves. Our results are summarized in Section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this article, we have performed a thorough linear and nonlinear analysis, from first principles, of electron acoustic excitations occurring in a nonthermal plasma consisting of hot $\\kappa$-distributed electrons, adiabatic cool electrons, and immobile ions.  First, we have derived a linear dispersion relation, and investigated the dependence of the dispersion characteristics on the plasma environment (degree of `suprathermality' through the parameter $\\kappa$, plasma composition, and thermal effects).  Then, we have employed the Sagdeev pseudopotential method to investigate large amplitude localized nonlinear electrostatic structures (solitary waves), and to determine the region in parameter space where stationary profile solutions may exist. Only negative potential solitons were found. The existence domain for solitons was shown to become narrower in the range of solitary wave speed, with an increase in the excess of suprathermal electrons in the hot electron distribution (stronger `suprathermality', lower $\\kappa$ value). The dependence of the soliton characteristics on the hot electron number density (through the parameter $\\beta$) and on the hot-to-cool electron temperature ratio $\\sigma$, were also studied. A series of appropriate examples of pseudopotential curves and soliton profiles were computed numerically, in order to confirm the predictions arising from the study of existence domains.  It may be added that ionic motion/inertia, here neglected, may also be included for a more accurate description, but is likely to have only minor quantitative effects.  We note, for completeness, that very recently two related papers have appeared with a scope apparently similar to that of the present article, viz., Refs. \\cite{Younsi2010, Sahu2010}. A word of comparison may therefore be appropriate here, for clarity. The latter authors~\\cite{Younsi2010} indicate that they are using a form of kappa distribution from one of the pioneering papers in the field.~\\cite{Thorne1991} Unfortunately, their expression for the characteristic speed $\\theta$ does not agree with the standard expression~\\cite{Thorne1991}, and thus the hot electron density does not take the usual form, Eq. (5).~\\cite{Baluku2008} Further, they use values of $\\kappa\\geq0.6$, i.e., well below the standard $\\kappa$-distribution cut-off of 3/2. Hence their results do not apply to the standard form of $\\kappa$ distribution.~\\cite{Hellberg2009}  As regards the other paper,~\\cite{Sahu2010} it did not consider existence domains at all (apart from an indirect mention of an upper limit in $M$, as commented on in Section \\ref{existence} above). Further, no account is taken in the paper of the possible effects of Landau damping on sustainable nonlinear structures, and a number of the figures relate to values of $\\beta$ (called $\\alpha$ in the paper) which lie in the unphysical, damped range.  We note that Ref.~\\onlinecite{Sahu2010} has also carried out a ``small amplitude\" calculation yielding double layers. However, when evaluated numerically, these turn out to be well beyond the range of small amplitude, and that raises some doubts about the validity of the results (their Figures 8-10). Further, let us take together their Figures 2 and 8, and consider the case of $\\kappa=3$ and $\\alpha=0.2$ (i.e., our $\\beta$). The latter is, of course, a value for which the linear wave is likely to be strongly Landau damped. It appears from the figures that a soliton occurs at $M=1.1$ with amplitude $\\sim0.9$, while a double layer occurs for $M=2.0$ with amplitude $\\sim0.7$. This combination of data does not satisfy the analytically-proven requirement that $\\partial\\Psi/ \\partial M <0$% ,~\\cite{Verheest2010,Verheest2010a,Verheest2010b}. There is no obvious reason why that should be the case, and there thus seems to be an error in at least one of these two figures.  Finally, we point out that we have not sought double layers in our calculations. However, we would be surprised if they did occur, as they are usually found as the upper limit to a sequence of solitons for a polarity for which there is no other limit. In this case there clearly is an upper limit for negative solitary waves, arising from the constraint $F_{2}(M)=0$. It is in principle possible for a double layer to occur at a lower value of $M$, and be followed by larger amplitude solitons at higher $M$, until the upper fluid cutoff such as a sonic point or an infinite compression cutoff is reached.~\\cite{Baluku2010} However, such behaviour depends on the the Sagdeev potential having a fairly complicated shape, with subsidiary local maxima, and we have not observed these for this model.  We are not aware of any experimental studies with which these theoretical results may be directly compared. However, it has previously been shown that wave data may be used to obtain an estimate for $\\kappa$, thus acting as a diagnostic for the distribution function~\\cite{Hellberg2000,Vinas2005}.  Similarly, in this case there are a number of indicators amongst our results which experimenters may wish to consider when interpreting observations. Thus, for instance, a lower normalized phase velocity of the linear electron-acoustic wave than would be predicted by a Maxwellian model (see Fig.~\\ref{fig1}) could be used to evaluate $\\kappa$.  Secondly, from Fig.~\\ref{fig2} one sees that in low-$\\kappa$ plasmas the range of normalized soliton speeds is both narrower and of larger value than one would expect for a Maxwellian. Thus, if solitons are found with normalized speeds around $M \\simeq0.4$, these can be understood only by allowing for additional suprathermal electrons (lower $\\kappa$). Further, from Fig.~\\ref{fig2} it follows that Maxwellian electrons give rise to a cutoff in the density ratio $\\beta\\simeq1$, and hence solitons observed in such plasmas can only be explained in terms of lower $\\kappa$.  From Fig.~\\ref{fig6} we note that at fixed values of the normalized soliton speed, $M$, the amplitudes of the perturbations of the normalized potential, cool electron density and cool electron speed due to the solitary waves all increase with decreasing $\\kappa$. This is related to the increase of the true Mach number $M/M_{1}$ for smaller $\\kappa$, as the phase velocity $M_{1}$ is decreased. Thus larger disturbances are likely to be associated with increased suprathermality.  Finally, turning to Fig.~\\ref{fig8}, two effects are observed: At fixed true Mach number, $M/M_{1}$, the soliton amplitude decreases with decreasing $\\kappa$ (increasing suprathermality). Despite that, the maximum values of soliton amplitude is found to occur not for a Maxwellian, but for the relatively low-$\\kappa$ values of around 2.5-3. Thus, again, large observed amplitudes are likely to be associated with a low-$\\kappa$ plasma.  Hence, as shown above, these results could assist in the understanding of solitary waves observed in two-temperature space plasmas, which are often characterized by a suprathermal electron distribution."
    },
    "1107/1107.0403_arXiv.txt": {
        "abstract": "We investigate the $Q$-ball decay in the gauge-mediated SUSY breaking. $Q$ balls decay mainly into nucleons, and partially into gravitinos, while they are kinematically forbidden to decay into sparticles which would be cosmologically harmful. This is achieved by the $Q$-ball charge small enough to be unstable for the decay, and large enough to be protected kinematically from unwanted decay channel. We can then have right amounts of the baryon asymmetry and the dark matter of the universe, evading any  astrophysical and cosmological observational constraints such as the big bang nucleosynthesis, which has not been treated properly in the literatures. ",
        "introduction": "Production of dark matter and baryon number of the universe is very important in cosmology. In the gauge-mediated supersymmetry (SUSY) breaking, the gravitino is natural candidate of the dark matter, while the Afflick-Dine (AD) baryogenesis \\cite{AD} is promising mechanism because thermal leptogenesis confronts serious cosmological gravitino problem, for example. It is well known that the AD condensate transforms into nontopological solitons, $Q$ balls, during the course of baryogenesis \\cite{KuSh,EnMc,KK1,KK2,KK3}. The charge of the $Q$ ball in the gauge mediation could be large enough to be stable against the decay into nucleons, and the $Q$ ball itself could be the dark matter of the universe \\cite{KuSh}. In this case, however, nondetections in direct and/or indirect $Q$-ball searches tell us that it may be difficult to simultaneously obtain enough baryon asymmetry of the universe\\footnote{% More recent observational limits tighter than those applied in Ref.\\cite{KK3} may wipe out the allowed regions.}.  On the other hand, the small charge $Q$ ball will decay into nucleons. At the same time, it could decay into gravitinos, which would have enough abundance to be the dark matter. This idea was considered in Refs.\\cite{ShKu,DoMc}\\footnote{ Its possibility was mentioned in Ref.\\cite{KK3}.}. In the former, the authors considered relatively small charge $Q$ balls which evaporate in the thermal bath, and imagined that the gravitino is directly produced from the $Q$ balls. In this case, however, most of the charge decays (evaporates) at the temperature around the electroweak scale, and sparticles are created, finally leading to the production of the next-to-lightest SUSY particle (NLSP), which thermalizes immediately. The gravitino could be produced by the NLSP decay, but the abundance of the gravitino is too small to be the dark matter of the universe, and the NLSP decay would destroy light elements during the big bang nucleosynthesis (BBN).  In the latter reference, the formed $Q$ ball is the gravity-mediation type, which is able to decay into NLSPs nonthermally. The NLSP decays into the gravitiono in the end, but they did not discuss the fact that the NLSP abundance is limited severely by the BBN constraints for $m_{3/2} \\gtrsim 1$ GeV,  so that the amount of the produced gravitino should be much smaller than that to be the dark matter \\cite{BBN} \\footnote{% Their model may work for a right-handed sneutrino LSP or if there is a small amount of R-parity violation\\cite{McDonald}. }.  In this article, we investigate a very simple scenario that the unstable $Q$ balls decay mainly into nucleons, partially into gravitinos with small branching ratio, while kinematically forbidden to NLSPs almost throughout the decay process, in the gauge-mediated SUSY breaking. The crucial observation reveals that the right amounts of the baryon asymmetry and the dark matter of the universe can be created through this process. The scenario is achieved by a couple of key ingredients. One is that the $Q$ ball decays from its surface \\cite{evap} so that the decay rate through the main channel has an upper bound, which is called the saturated rate, while the decay into gravitino is not saturated. Another is that a large enough $Q$ ball forbids the decay into the NLSP, while allowing the decay into nucleons and gravitinos, having an appropriate lifetime to decay at the temperature of $O$(10 MeV).  The whole scenario proceeds as follows. The AD field starts the oscillation (rotation) after inflation. The rotation has an oblate orbit in the potential. Due to the negative pressure, the AD field feels spatial instabilities, which grow into $Q$ and anti-$Q$ balls. The $Q$ and anti-$Q$ balls decay into nucleons and anti-nucleons to create the baryon asymmetry in the universe, while producing small amount of gravitinos which become the dark matter of the universe. NLSPs would be produced only at the very end of the decay process when the charge of the Q ball becomes small enough to open the kinematically allowed channel to the NLSP production\\cite{KT}, and its abundance is small enough to evade any cosmological disasters such as spoiling the successful BBN.  The structure of the article is as follows. In the next section, we show the features of the $Q$ ball in the gauge-mediated SUSY breaking model, and the details of the decay process of the $Q$ ball are described in Sec.~\\ref{Qdecay}. In Sec.\\ref{abundances}, we estimate the abundances of the baryon number, the dark matter, and the NLSP. We impose BBN constraints on the NLSP abundances to derive the allowed gravitino mass range in Sec.\\ref{BBNconstraints}. In Sec.\\ref{scenario}, we show the region in the model parameter space where the successful scenario exists. Section \\ref{conclusion} is devoted to our conclusions. ",
        "conclusions": "\\label{conclusion} We have investigated the scenario of simultaneous production of the baryon asymmetry and the dark matter of the universe through the $Q$-ball decay in the gauge-mediated SUSY breaking model. This is simply achieved by the $Q$ ball with charge small enough to decay into nucleons to create baryon number of the universe. We have calculated the branching ratio of the gravitino to find that it is small. Therefore, in order to obtain the observed baryon to dark matter abundance ratio, the AD field should rotate in the oblate orbit so that both $Q$ and anti-$Q$ balls are created to suppress the baryon number compared to the gravitino dark matter.  At the same time, the charge should be large enough to kinematically forbid the decay into NLSPs. This evades the serious BBN constraints on the NLSP abundance, which was not properly taken into account in the literatures. NLSPs are actually produced at the very end of the $Q$-ball decay when the charge becomes small enough to open the corresponding decay channel.  Including all these considerations, we have used the BBN constraints for the bino, stau, and sneutrino NLSPs in Ref.\\cite{BBN}, to get the allowed range of the gravitiono mass. We have found that typically $m_{3/2} \\lesssim 1 \\ (10^{-2})$ GeV for $m_{\\rm NLSP}=300 \\ (100)$ GeV.  We have also shown the allowed region in $Q-M_F$ plane for $m_{\\rm NLSP}=300$ GeV. We thus have found the production of right amounts of both the baryon asymmetry and the gravitino dark matter of the universe is successful when the $Q$ ball has the charge $Q \\simeq 10^{23}$ for $M_F \\simeq 10^6 - 10^7$ GeV and $m_{3/2} \\simeq 10^{-2} - 10^{-1}$ GeV. It is naturally achieved, for example, for the $n=6$ $udd$ direction. In this case, the nonrenormalizable superpotential $W_{\\rm NR}= (udd)^2/M_{\\rm P}^3$ determines the field amplitude at  the onset of the oscillation $\\phi_{\\rm osc} \\simeq (M_F^2M_{\\rm P}^3)^{1/5}$, and also leads to the $A$-term of the form $V_A \\sim a m_{3/2} W_{\\rm NR} + {\\rm h.c.}$, and the scenario is successful for $M_F\\sim10^7$ GeV, $m_{3/2}\\sim 0.01$ GeV, and $a \\sim 0.1$, which leads to $\\phi_{\\rm osc} \\sim 7\\times 10^{13}$ GeV, and $Q \\sim 10^{23}$."
    },
    "1107/1107.0918_arXiv.txt": {
        "abstract": "{ As part of our radial velocity planet-search survey performed with SARG at TNG, we monitored the components of HD 132563 for ten years. It is a binary system formed by two rather similar solar type stars with a projected separation of 4.1 arcsec, which corresponds to 400 AU at the distance of 96 pc. The two components are moderately metal-poor ([Fe/H]=--0.19), and the age of the system is about 5 Gyr. We detected RV variations of HD 132563B with period of 1544 days and semi-amplitude of 26 m/s. From the star characteristics and line profile measurements, we infer their Keplerian origin. Therefore HD 132563B turns out to host a planet with a projected mass $m \\sin i=1.49~M_{J}$ at 2.6 AU with a moderately eccentric orbit ($e=0.22$). The planet around HD 132563B is one of the few that are known in triple stellar systems, as we found that the primary HD 132563A is itself a spectroscopic binary with a period longer than 15 years and an eccentricity higher than 0.65. The spectroscopic component was not detected in adaptive-optics images taken with AdOpt@TNG, since it expected at a projected separation that was smaller than 0.2 arcsec at the time of our observations. A small excess in K band difference between the components with respect to the difference in V band is compatible with a companion of about $0.55~M{_\\odot}$. A preliminary statistical analysis of when planets occur in triple systems indicate a similar frequency of planets around the isolated component in a triple system, components of wide binaries and single stars. There is no significant iron abundance difference between the components ($\\Delta$[Fe/H]=$0.012\\pm 0.013$ dex). The lack of stars in binary systems and open clusters showing strong enhancements of iron abundance, which are comparable to the typical metallicity difference between stars with and without giant planets, agrees with the idea that accretion of planetary material producing iron abundance anomalies over 0.1 dex is rare. } ",
        "introduction": "\\label{s:intro}  Planets in multiple stellar systems play a special role in determinations of the frequency of planets in the Galaxy, because more than half of solar type stars are in multiple systems (Duquennoy \\& Mayor \\cite{duq}). Furthermore, planets in binaries allow one to better understand the processes of planet formation and evolution, thanks to the dynamical effects induced by the stellar companions on the protoplanetary disks and on an already formed planetary system (Marzari et al.~\\cite{marzari10}). Finally, since the frequency of (wide) binaries seems related to the density of the formation environment (Lada \\& Lada \\cite{lada03}; Goodwin \\cite{goodwin10}) some insight into the impact of this on the planet formation mechanism can be derived (Malmberg et al.~\\cite{malmberg07}).  In spite of selection criteria against binaries in many planet-search surveys, several tens of planets are currently known in multiple systems (Desidera \\& Barbieri \\cite{db07}, Mugrauer \\& Neuhauser \\cite{mugrauer09}). A variety of binary configurations are observed: several of the planets are in rather wide binaries, for which the effects caused by the companion on the planetary system are expected to be small. However, there are few planet hosts with companions at a separation that is small enough (e.g. Hatzes et al.~\\cite{hatzes03}; Chauvin et al.~\\cite{chauvin11}) to make it difficult to allow planet formation according to our current knowledge (see e.g. Thebault \\cite{thebault11}) even if they are on dynamically stable orbits. The variety of system configurations is further extended by the discovery of planets in triple systems and, more recently, of planets in circumbinary orbits around close eclipsing binary systems (e.g Lee et al.~\\cite{lee09}), when using the timing technique.  Planet properties also appear to be affected by the presence of companions, as is the case for the mass distribution, with massive planets in close orbits found primarily in binaries (Zucker \\& Mazeh \\cite{zucker02}; Udry et al.~\\cite{udry03}; Desidera \\& Barbieri \\cite{db07}), and for planet eccentricity, which is enhanced in binary systems, possibly through the Kozai resonance (Tamuz et al.~\\cite{tamuz08}). The circumbinary planets found so far around eclipsing binaries are all rather massive and have long periods. The possible link of these features with the original system configuration and its evolution (all the central binaries include a remnant object) needs to be investigated.  The general radial velocity (RV) surveys include several binaries in their samples, but in most cases only the brightest star in the system is under RV monitoring. Dedicated surveys of planets in binary systems with well-defined properties help for better understanding the impact of binarity on planet formation. Some of these projects are focused on spectroscopic binaries (Eggenberger \\cite{egg09}; Konacki et al.~\\cite{konacki}), while others study visual binaries using RV (our survey,  Desidera et al.~\\cite{sargbook}; Toyota et al.~\\cite{toyota09}) and astrometry (Muterspaugh et al.~\\cite{muterspaugh10}). Our ongoing survey using SARG at TNG is the first planet-search survey specifically devoted to binary systems. It targets moderately wide binaries (typical separation 100-500 AU) with similar components (twins). The sample includes about 50 pairs of main sequence solar-type stars with projected separation larger than 2 arcsec, distance smaller than 100 pc and magnitude difference $\\Delta V < 1$ mag. The sample selection and the survey itself are described in Desidera et al.~(\\cite{sargbook}).  Unlike most of the multiple systems included in the RV surveys, both components of the pair are systematically observed in our survey, usually with the same time sampling. This allows the investigation of possible factors driving planets in binary systems to occur and of possible effects induced, directly or indirectly, by the presence of planets, such as altering chemical abundances by accreting metal-rich planetary material (Desidera et al.~\\cite{chem2}).  In this paper we present the interesting case of the \\object{HD 132563} system, whose components show both RV variations: those with amplitudes larger than 1 km/s for the primary, implying most likely a stellar companion, and those with amplitudes of a few tens of m/s for the secondary, suggesting there is a planetary companion. The paper is organized as follows. In Sect.~\\ref{s:obs} we present the observations taken at TNG using SARG (high-resolution spectroscopy) and AdOpt (direct imaging). In Sect.~\\ref{s:star} we present the properties of the components. In Sect.~\\ref{s:hd132563b} we present the evidence for a planetary companion around \\object{HD132563B}. In Sect.~\\ref{s:hd132563a} we show that \\object{HD132563A} is a long-period spectroscopic binary, and we discuss the clues to the mass and the orbital parameters of the companion. In Sect.~\\ref{s:binorbit} we present the clues to the orbital parameters of the wide pair and its possible influence on the planet orbiting HD 132563B. In Sect.~\\ref{s:discussion} we discuss the results, with emphasis on planets in triple systems and on the metallicity difference between components with and without planets. In Sect.~\\ref{s:conclusion} we summarize our conclusions. ",
        "conclusions": "\\label{s:conclusion}  We have presented the discovery of a giant planet with a projected mass $m \\sin i = 1.49~M_{J}$ at 2.6 AU in a moderately eccentric orbit around HD132563B. This is the first planet detected using SARG at TNG as part of the survey that monitored about 100 stars that are members of wide binaries with similar components.  The star has a wide companion (HD132563A) at a projected separation of 400 AU, which we found to be  a spectroscopic binary with a period of about 25-30 yr in a highly eccentric orbit. Taking various observational constraints into account (unresolved photometry of HD132563A, line profile analysis, adaptive optics imaging, astrometric data), we concluded that most likely the companion has a mass of about $0.5-0.6~M_{\\odot}$. Therefore, HD 132563Bb is one of eight planets known to be in stellar triple systems. A tentative statistical analysis indicates that the isolated component in a hierarchical triple system is not hostile to planet formation, with planet frequency as high as for single stars and the components of wide binaries.  HD132563B is a slightly metal-poor star with a mass very close to that of Sun. Its companion is less than 200~K warmer. Such a small temperature difference, coupled with the small convective envelope of the star, makes this system ideal for exploring accretion of metal-rich material of planetary origin. The differential abundance analysis (already performed in Desidera et al.~\\cite{chem2}) yielded no significant abundance difference, with an error small enough to exclude accretion of about $5~M_{\\oplus}$ of meteoritic material, which is the quantity that should have been accreted by the Sun during its main sequence lifetime.  This result, along with the lack of objects in binary systems and open clusters showing large enhancements of iron abundance that are linked to the evolution of planetary systems, agrees with the idea that large abundance anomalies, comparable to the typical metallicity difference between stars with and without giant planets, are rare."
    },
    "1107/1107.4827_arXiv.txt": {
        "abstract": "A pixel analysis is carried out on the interacting face-on spiral galaxy NGC 5194 (M51A), using the {\\it HST}/ACS images in the F435W, F555W and F814W ($BVI$) bands. After $4\\times4$ binning of the {\\it HST}/ACS images to secure a sufficient signal-to-noise ratio for each pixel, we derive several quantities describing the pixel color-magnitude diagram (pCMD) of NGC 5194: blue/red color cut, red pixel sequence parameters, blue pixel sequence parameters and blue-to-red pixel ratio. The red sequence pixels are mostly older than 1 Gyr, while the blue sequence pixels are mostly younger than 1 Gyr, in their luminosity-weighted mean stellar ages. The color variation in the red pixel sequence from $V=20$ mag arcsec$^{-2}$ to $V=17$ mag arcsec$^{-2}$ corresponds to a metallicity variation of $\\Delta$[Fe/H] $\\sim2$ or an optical depth variation of $\\Delta\\tau_V\\sim4$ by dust, but the actual sequence is thought to originate from the combination of those two effects. At $V<20$ mag arcsec$^{-2}$, the color variation in the blue pixel sequence corresponds to an age variation from 5 Myr to 300 Myr under the assumption of solar metallicity and $\\tau_V=1$. To investigate the spatial distributions of stellar populations, we divide pixel stellar populations using the pixel color-color diagram and population synthesis models. As a result, we find that the pixel population distributions across the spiral arms agree with a compressing process by spiral density waves: dense dust $\\rightarrow$ newly-formed stars. The tidal interaction between NGC 5194 and NGC 5195 appears to enhance the star formation at the tidal bridge connecting the two galaxies. We find that the pixels corresponding to the central active galactic nucleus (AGN) area of NGC 5194 show a tight sequence at the bright-end of the pCMD, which are in the region of $R\\sim100$ pc and may be a photometric indicator of AGN properties. ",
        "introduction": "NGC 5194 is one of the nearest galaxies with grand design spiral arms, also known as M51A\\footnote{The name `M51' often designates the interacting system of NGC 5194 + 5195, while it sometimes indicates NGC 5194 only. In this paper, to avoid the terminological confusion, we use the name `M51' only when we mean the NGC 5194 + 5195 system. } or the Whirlpool galaxy. NGC 5194 is a very interesting object not only because we can inspect its spiral arm structures with close and face-on views, but also because NGC 5194 is known to be interacting with its companion galaxy NGC 5195. The morphological type of NGC 5194 is Sbc \\citep{dev91} and that of NGC 5195 is SB0 \\citep{san87}. Due to such interesting properties, the entire area of the interacting system NGC 5194 + 5195 was covered using the {\\it Hubble Space Telescope} ({\\it HST}) and Advanced Camera for Surveys (ACS) with very high resolution by the Hubble Heritage Team \\citep{mut05}, despite its large angular size ($\\sim9'$).  Owing to its scientific attraction and high-quality data, NGC 5194 has been investigated by many researchers. One of the active studies on NGC 5194 using the {\\it HST} data is the search and investigation of star clusters in it. Until now, about 2000 star clusters in M51 have been found and their physical properties, such as age and mass, were derived, revealing that the star cluster formation rate increased significantly during the period of $100-250$ Myr ago, which is consistent with the epoch of the dynamical encounters of the two galaxies \\citep[e.g.][]{lee05,hwa08,hwa10}. Bright stars are resolved as point sources in NGC 5194, like the star clusters. \\citet{kal10} presented a photometric study of the stellar associations across the entire disk of NGC 5194, determining the age, size, mass, metallicity and dust content of each association as a function of location. In addition to the stellar properties, the collective properties of the H{\\scriptsize II} regions in M51 were also investigated using the \\emph{HST} $H{\\alpha}$-band data in \\citet{lee11}. They found that the $H{\\alpha}$ luminosity functions have different slopes for different parts in M51, which implies that the H{\\scriptsize II} regions in the interarm region are relatively older than those in the other parts of M51.  The effect of interaction on the properties of NGC 5194 is a topic studied by many researchers. \\citet{dur03} investigated the kinematics of planetary nebulae in M51's tidal debris, finding that the tidal debris consists of two discrete structures: NGC 5195's own tidal tail and diffuse material stripped from NGC 5194, which is consistent with the `single-passage' model for the encounter with a 2:1 mass ratio. \\citet{mei08} tested whether the strong spiral structure is transient (i.e., interaction-driven) or steady, based on the radial dependence of the pattern speed, showing two distinct pattern speeds inside 4 kpc at their derived major axis, which are indicative of a long-lasting spiral structure.  The studies on the nuclear region of NGC 5194 are also active. \\citet{gri97} analyzed the central region of NGC 5194 imaged with the {\\it HST}, finding a complicated distribution of dust lanes. Those dust lanes are roughly aligned with the major axis of the bar and may be transporting gas to the AGN in the nucleus. Later, a Chandra observation on the AGN and nuclear outflows of NGC 5194 was carried out \\citep{ter01} and showed a low-luminosity Seyfert 2 nucleus, southern extra-nuclear cloud and northern loop. The morphology of the extended emission is very similar to radio continuum and optical emission-line images. \\citet{lam02} inspected the ongoing massive star formation in the bulge of NGC 5194 using the {\\it HST}/WFPC2 observation data and revealed that the bulge around the nucleus is dominated by a smooth reddish background population with overimposed dust lanes. They found 30 bright point-like sources in the bulge within 110-350 pc of the nucleus, which are probably isolated massive stars. \\citet{mat04,mat07} studied the jet-disturbed molecular gas near the NGC 5194 AGN, which showed the presence of a dense circumnuclear-rotating disk, which may be a reservoir that fuels the active nucleus and obscures it from direct optical view.  Those various studies revealed many interesting properties of NGC 5194, but there remains a useful analysis method not previously tried for NGC 5194: {\\it pixel analysis}. Recently, pixel analysis methods started to be used for investigation of various galaxies, thanks to the observational facilities and techniques that provide high spatial resolution and high signal-to-noise ratio (S/N) even for individual pixels. One of the well-known pixel analysis methods is the {\\it pixel-z} technique established by \\citet{con03} to investigate the star formation history of galaxies measured from individual pixels. \\citet{joh05} carried out pixel-by-pixel SED fitting for a galaxy hosting a compact steep-spectrum radio source MRC B1221$-$423, based on $BVRIK$ imaging and spectroscopic observations. More recently, the pixel-z technique was applied for 44,964 galaxies to derive their ages, star formation rates, dust obscurations, metallicities, and morphology -- density relations \\citep{wel08,wel09}. \\citet{wij10} extracted the physical properties of the SDSS galaxies with the pixel-z method, relating the spatial distribution of the physical properties with the morphological properties of galaxies.  The pixel-z method is powerful, but it is based on SED fitting, which depends on SED templates and fitting algorithms and requires images in many (at least four) bands. A simpler way is the pixel color-magnitude or color-color diagram analysis, which was tried a long time ago for estimating stellar populations in unresolved galaxies \\citep[e.g.][]{bot86}. \\citet{kas03} studied the stellar populations of merging galaxies NGC 4038/4039 in the $BVK$ bands, deriving pixel-by-pixel maps of the distributions of stellar populations and dust extinction. From the pixel color-color diagrams (pCCDs) of those galaxies and population synthesis models, they estimated the approximate age and reddening of each pixel in those galaxies. Similarly, \\citet{deg03} analyzed the pixel color-magnitude diagram (pCMD) and pCCD, to study the star formation histories of the Mice (NGC 4676) and Tadpole (UGC 10214) interacting galaxies. The spatial distributions of pixels in different domains of the pCMD and pCCD were investigated and revealed that galaxy interactions significantly affect stellar populations.  More recently, a systematic comparison of the pCMD features between galaxies with different morphologies was carried out by \\citet{lan07} for a large number of galaxies. They studied 69 nearby galaxies chosen to span a wide range of Hubble types using their pCMDs and found that early-type galaxies exhibit clear \\emph{prime sequences} but some of them have dust-originated \\emph{red hook} features. They also found that lenticular galaxies show pCMD features similar to those of elliptical galaxies but tend to have larger color dispersions and that some pCMDs of face-on spirals show \\emph{inverse-L} features. In addition, they developed quantitative methods to characterize the pCMDs, such as the blue-to-red light ratio and total color dispersion. \\citet{lan07} suggested a leading concept in the methodology of pCMD analysis: the \\emph{quantification} of pCMD features. By quantifying the various features in pCMDs, it is possible to compare the pCMDs systematically between different galaxies, which is a totally new methodology in galaxy studies using simple (but highly-resolved) image data. However, since the quantified parameters in \\citet{lan07} are so simple they are not enough to describe the various features of pCMDs, that is, the pCMD analysis method is not perfectly established yet.  In this paper, we present a pixel analysis of the interacting face-on spiral galaxy, NGC 5194. The main goals of this paper are 1) to improve the understanding of NGC 5194 properties using pixel analysis methods and 2) to establish the pCMD analysis methods for future photometric studies of galaxies. For the second goal, NGC 5194 (or M51) is an ideal target, because it is a very nearby face-on galaxy, and the entire area was covered by the {\\it HST} with very high spatial resolution. With the high-resolution image data, how the pCMD features vary as a function of spatial resolution can be tested.  The outline of this paper is as follows. Section 2 describes the data and \\S3 explains the pixel binning and analysis area. The results of the pixel analysis are shown in \\S4 and discussed in \\S5. Section 6 gives the conclusion. Throughout this paper, we adopt the cosmological parameters: $h=0.7$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{M}=0.3$. ",
        "conclusions": "We carried out a pixel analysis of the interacting face-on spiral galaxy NGC 5194, using {\\it HST}/ACS images in the $BVI$ bands. Three major results were presented. First, we derived several quantities describing the pixel color-magnitude diagram (pCMD) features of NGC 5194, such as blue/red color cut, red pixel sequence parameters, blue pixel sequence parameters and blue-to-red pixel ratio. Those parameters reflect the internal properties of NGC 5194, such as age, metallicity, dust content and galaxy morphology. In the luminosity-weighted mean stellar age estimation with solar metallicity and $\\tau_V=1$ assumptions, the red sequence pixels are mostly older than 1 Gyr, while the blue sequence pixels are mostly younger than 1 Gyr. The color variation in the red pixel sequence from $V=20$ mag arcsec$^{-2}$ to $V=17$ mag arcsec$^{-2}$ corresponds to a metallicity variation of $\\Delta$[Fe/H] $\\sim2$ or an optical depth variation of $\\Delta\\tau_V\\sim4$ by dust, but the actual sequence is thought to originate from the combination of those two effects. At $V<20$ mag arcsec$^{-2}$, the color variation in the blue pixel sequence corresponds to age variation from 5 Myr to 300 Myr in solar metallicity and $\\tau_V=1$ assumptions. The spatial resolution dependence of the pCMD features was tested for the future comparisons between galaxies with different image resolutions, and we found that the pixel resolution needs to be better than 100 pc pixel$^{-1}$ for the investigation of pCMD features.  Second, we found that the spatial distributions of pixel stellar populations are significantly affected by the spiral arm patterns and the tidal interaction with NGC 5195. The pixel population distributions across the spiral arms agree with a compressing process by spiral density waves: dense dust (before-arm) $\\rightarrow$ newly-formed stars (after-arm). The interaction between NGC 5194 and NGC 5195 also enhances the star formation at the tidal bridge area connecting the two galaxies.  Finally, we found that the pixels corresponding to the central area (around the AGN) of NGC 5194 show a tight sequence in the bright-end of the pCMD, of which spatial extent is $R\\sim100$ parsecs. The origin and mechanism of the tight bright-end pixel sequence are not sufficiently explained in this paper, but the sequence seems to be worth a comparison between various AGNs using high-resolution imaging data in the future, because it may be a photometric indicator of AGN properties."
    },
    "1107/1107.2147.txt": {
        "abstract": "We present uncontaminated rest-frame $u-R$ colors of 78 X-ray-selected AGN hosts at $0.5 < z < 1.5$ in the Chandra Deep Fields measured with HST/ACS/NICMOS and VLT/ISAAC imaging.  We also present spatially-resolved $NUV-R$ color gradients for a subsample of AGN hosts imaged by HST/WFC3.  Integrated, uncorrected photometry is not reliable for comparing the mean properties of soft and hard AGN host galaxies at $z \\sim 1$ due to color contamination from point-source AGN emission.  We use a cloning simulation to develop a calibration between concentration and this color contamination and use this to correct host galaxy colors.  The mean $u - R$ color of the unobscured/soft hosts beyond $\\sim6$ kpc is statistically equivalent to that of the obscured/hard hosts (the soft sources are $0.09 \\pm 0.16$ magnitudes bluer).  Furthermore, the rest-frame $V-J$ colors of the obscured and unobscured hosts beyond $\\sim6$ kpc are statistically equivalent, suggesting that the two populations have similar distributions of dust extinction.  For the WFC3/IR sample, the mean $NUV - R$ color gradients of unobscured and obscured sources differ by less than $\\sim 0.5$ magnitudes for $r > 1.1$ kpc.  These three observations imply that AGN obscuration is uncorrelated with the star formation rate beyond $\\sim1$ kpc.  These observations favor a unification scenario for intermediate-luminosity AGNs in which obscuration is determined geometrically. Scenarios in which the majority of intermediate-luminosity AGN at $z\\sim1$ are undergoing rapid, galaxy-wide quenching due to AGN-driven feedback processes are disfavored. ",
        "introduction": "The physical connection between embedded supermassive black holes and their host galaxies remains an active topic of research.  The discovery of correlations between bulge properties and black hole mass \\citep[e.g., stellar mass and velocity dispersion,][]{mag98, fer00, geb00} imply that this connection spans an astonishing range of spatial scales;  the AU-sized accretion disk whose characteristics determine the black hole growth rate must receive information about the conditions of the galaxy on scales of kiloparsecs, or eight orders of magnitude larger.  These correlations between black hole mass and various bulge properties hint at simultaneous evolution of AGNs and their host galaxies.  The processes that are involved in the maintenance of the $M-\\sigma$ relation may also govern the evolution of host galaxy morphology and color.  In the local universe, galaxies exhibit bimodal distributions of several critical parameters, including rest-frame optical color, star formation rate, and gas fraction \\citep[][etc]{str01, kau03, bal04}.  Star formation must be suppressed significantly to explain the low star formation rates and the redness of early-type, passively evolving galaxies today.  There is circumstantial evidence for the participation of AGN in this suppression of star formation, or ``quenching,'' including their presence in the green valley of the galaxy color-magnitude diagram (CMD), intermediate between the blue cloud and red sequence \\citep{kau03, nan07, sch09, car10}, although the AGN fraction does not appear to be higher in the green valley \\citep{xue10}.  Major mergers of gas-rich galaxies can result in the rapid growth of central black holes.  In this picture, the collision of gas clouds saps angular momentum, efficiently sending gas to the central regions of the galaxy \\citep{her89}, where it is largely consumed in star formation and partially accreted onto one or more BHs.  In several models, the ``M-$\\sigma$'' relation is maintained through self-regulating processes involved in galaxy and supermassive black hole (SMBH) growth \\citep[][and references within]{hop05, dim05, hop09}.    This self-regulation may be modulated by star formation, in which momentum injection by Type Ia supernovae and radiation pressure from O stars disrupt further star formation \\citep{mur05}.  However, the immense energies associated with gravitational accretion onto a SMBH dominate the energy budget of the galaxy; this suggests that the AGN itself may provide sufficient energy to remove gas and limit the growth of the BH \\citep{sil98, hae98, kin03, dim05, hop05, spr05, hop06, men08}.   A further consequence of the most energetic AGN feedback processes may be the expulsion of a majority of the galactic gas \\citep{san88, dim05, hop06}.  This quenching of star formation may be a mechanism by which AGN activity controls the evolution of galaxies from the blue cloud to the red sequence.  One feature of a model that includes rapid, AGN-induced quenching is an evolution between obscured and unobscured AGN states \\citep{hop06, men08}.  The ejection of gas/dust in the central core of the galaxy is quickly followed by gas removal on kpc-scales.  This process naturally produces a positive correlation between AGN obscuration and the mass fraction of young stars (a consequence of sufficient gas supply) in the host galaxies \\citep{men08}.  Unobscured AGN are expected to reside in relatively gas-free host galaxies, cleared by a previous blowout phase.  These hosts are expected to be redder than the hosts of obscured AGN, which would have higher star-formation rates.  It is not clear how this theoretical evolution between obscured and unobscured states is consistent with the observed AGN unification paradigm.  AGN unification schemes attempt to explain the properties of the many observed types of AGN with a common physical model.  The most successful unification picture postulates that the differences between Seyfert types in the local universe are due to variation of the orientation of the observer relative to a centralized dusty torus \\citep{ant93, urr95}.  Since obscuration is solely determined by geometry in a unified scheme, there is no reason to expect differences between the properties of the host galaxies of obscured and unobscured AGN.  Many comparisons of obscured and unobscured AGNs have been performed at low and high redshift, using both optical- and X-ray-based diagnostics of AGN obscuration and various measures of star formation rate or gas content.  Locally, obscured Seyferts are associated with significant starforming activity, in some cases more than comparable unobscured hosts \\citep{cid95, gon01}.  \\citet{lac07} finds that obscured Type II QSOs at redshifts of $0.3 < z < 0.8$ are associated with significant star forming activity.  However, using HI measurements, unobscured broad-line AGN hosts at low-redshift are observed to be gas-rich \\citep{ho08}.  \\citet{sch09} find no dependence of host galaxy color on obscuring column for a small sample of SWIFT BAT AGN \\citep{geh04, tue08} in the local universe.  \\citet{pag04} find stronger submillimeter emission in an X-ray obscured sample compared to an X-ray unobscured sample, implying a higher star formation rate in the obscured sample.  \\citet{gil09} find that the clustering lengths of $z \\sim 1$ X-ray selected AGNs with broad-lines and those that lack broad lines are similar.  They also note that the clustering length of the X-ray hard sources is statistically similar to that of the soft sources, although the errors are dominated by number statistics.  \\citet{sha10} measure the star formation rates of X-ray selected AGNs in GOODS-North with Herschel far-infrared fluxes, finding no correlation with X-ray absorbing column.  \\citet{pie10b}, CP10 hereafter, uses optical colors of X-ray selected AGNs at $z \\sim 0.7$ in AEGIS to compare hard and soft sources.  Using colors measured with both integrated and extended apertures, they find that unobscured sources are bluer than obscured sources in rest-frame $U-V.$  \\subsection{Comparing the Colors and Color Gradients of Obscured and Unobscured AGN}  Generally, understanding the relationship between AGN host galaxy properties on kpc-scales and the properties of the AGN on smaller scales can be helpful in constraining the role of AGN in the evolution of galaxies.  We focus on comparing the host properties of obscured and unobscured sources, given the hypothesis that unobscured sources will reside in redder hosts in a feedback scenario involving AGN-driven blowouts.  We adopt a two-pronged approach:  We utilize two imaging datasets of different spatial resolution, allowing us to probe host properties on two spatial scales.  We first measure the rest-frame $u-R$ colors of a large sample of AGN host galaxies at high galactic radii (r $\\sim6-12$ kpc) and search for correlations of color with AGN obscuration or hardness ratio.  The combination of optical and infrared colors is a powerful diagnostic of the presence of young stars \\citep{gil02}, particularly rest-frame $u-R$, or observed $V-J$ at $z\\sim0.8$.  This sample is limited in spatial resolution by the infrared imaging data ($0\\farcs3$ - $0\\farcs6$ for ISAAC imaging and $\\sim 0\\farcs25$ for NIC3).  We then turn to a smaller sample of AGN host galaxies imaged by the Wide Field Camera 3 Infrared Channel (WFC3/IR) on HST as part of the Early Release Science (ERS) demonstration program \\citep{win10}.  The improved spatial resolution of $\\sim0\\farcs2$ enabled by WFC3 allows us to probe rest-frame $NUV-R$ color closer to the central regions of host galaxies at $r \\sim 1$ kpc, where AGN-induced winds may be the most energetic and color effects potentially more drastic.  Section 2 identifies the VLT/ISAAC and ACS data we use in the Great Observatories Origins Deep Survey (GOODS) field South \\citep{gia04} and the sample selection.  Section 3 details the methodology used to obtain uncontaminated aperture photometry of AGN host galaxies in these fields.  Section 4 describes a suite of cloning simulations used to check the analysis and derive error bars.  Section 5 presents results for all samples and Section 6 concludes.  We assume a flat cosmology throughout, with $H_0 = 70$ km/s/Mpc, $\\Omega_M = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": ""
    },
    "1107/1107.3546_arXiv.txt": {
        "abstract": "In models of dark matter (DM) with Sommerfeld-enhanced annihilation, where the annihilation rate scales as the inverse velocity, $N$-body simulations of DM structure formation suggest that the local annihilation signal may be dominated by small, dense, cold subhalos. This contrasts with the usual assumption of a signal originating from the smooth DM halo, with much higher velocity dispersion. Accounting for local substructure modifies the  parameter space for which Sommerfeld-enhanced annihilating DM can explain the PAMELA and \\emph{Fermi} excesses. Limits from the inner galaxy and the cosmic microwave background are weakened, without introducing new tension with substructure-dependent limits, such as from dwarf galaxies or isotropic gamma-ray studies. With substructure, previously excluded parameter regions with mediators of mass $\\sim$ 1-200 MeV are now easily allowed. For $\\mathcal{O}$(MeV) mediators, subhalos in a specific range of host halo masses may be evaporated, further suppressing diffuse signals without affecting substructure in the Milky Way. ",
        "introduction": "Interest in dark matter (DM) annihilation has been boosted in recent years as a consequence of a number of results from cosmic ray (CR) experiments. The PAMELA finding \\cite{Adriani:2008zr} of a rise in the positron fraction at high ($\\sim 10-100 \\gev$) energies coupled with harder than expected $e^+ + e^-$ spectra from \\emph{Fermi} \\cite{Abdo:2009zk, Ackermann:2010ij} and ATIC \\cite{aticlatest, Panov:2011zw} point to the existence of a new, primary source of high energy $e^+e^-$.  A more recent \\emph{Fermi} measurement of the positron fraction at $20-200 \\gev$ \\cite{FermiLAT:2011ab} confirms the rise reported by PAMELA.  Attempts to explain the new $e^+ + e^-$ component of cosmic rays as DM annihilation products are immediately confronted by several difficulties: the $e^+$ and $e^-$ fluxes implied by the high-energy excesses are $\\mathcal{O}(100-1000)$ times larger than expected from annihilation of a $\\sim$ TeV-mass thermal relic, no corresponding excess is seen in the antiproton signal, and the positron spectrum is too hard to arise from DM annihilation into weak gauge bosons or hadrons. Models of TeV-scale DM coupled to light (MeV-GeV) force carriers  \\cite{Finkbeiner:2007kk,ArkaniHamed:2008qn, Pospelov:2008jd} seek to address all three issues: kinematically forbidding the antiprotons, producing boosted positrons (giving rise to hard spectra), and allowing Sommerfeld enhancement of the annihilation cross section at low velocities \\footnote{The Sommerfeld enhancement was first discussed in the context of DM by \\cite{Hisano:2003ec, Hisano:2004ds}.}. Such models can explain the CR excesses while still yielding the appropriate thermal relic abundance, and without appealing to a local over-density of DM \\cite{Finkbeiner:2010sm}.  Most studies of such light-mediator models have assumed a smooth, Maxwellian halo in the local (within $\\sim 1$ kpc) neighborhood, where the positrons and electrons observed by PAMELA and \\emph{Fermi} are thought to originate.   But analytical arguments and $N$-body simulations both suggest a richer halo substructure, with a non-negligible fraction of DM residing in self-bound subhalos as small as $10^{-6} M_\\odot$, as well as tidal streams and other structures.  Neglecting these subhalos would at first appear to be a well-justified approximation;  $N$-body simulations suggest that they contribute only an $\\mathcal{O}(1)$ factor to the density-squared integral, and hence to the annihilation rate for conventional Weakly Interacting Massive Particles (WIMPs). However, bound subhalos typically have velocity dispersions much smaller than the $\\sim$150 km/s of the smooth halo.  As Sommerfeld enhancement confers an additional $1/v$-$1/v^2$ scaling to the annihilation cross section (down to some saturation velocity, below which $\\langle \\sigma v \\rangle$ is increased by a constant \\emph{saturated enhancement}), even $\\mathcal{O}(1)$ contributions to the density-squared integral can change the predicted Sommerfeld-enhanced DM annihilation rate by an order of magnitude or more.  As such, annihilation in local substructure can overwhelmingly dominate the smooth halo signal in models with Sommerfeld-enhanced annihilation (related effects have been studied in \\cite{Lattanzi:2008qa, Kuhlen:2009is, Kuhlen:2009kx, Bovy:2009zs,Yuan:2009bb, Vincent:2010kv, Kistler:2009xf, Kamionkowski:2010mi}).    We show in Fig. \\ref{fig:satlocalratio} the size of the saturated Sommerfeld enhancement relative to the Sommerfeld enhancement in the smooth halo. For heavier mediator masses, the difference is only a factor of a few (except on resonances), but for mediator masses $\\lesssim 100 \\mev$ the saturated enhancement is generically a factor of 10-100 times larger than the smooth-halo boost.  \\begin{figure*} \\includegraphics[width=0.45\\textwidth]{satlocalratio.pdf} \\caption{The ratio of the saturated Sommerfeld enhancement to the smooth halo enhancement as a function of force carrier mass, assuming a local 1D velocity dispersion of $\\sigma = 150$ km/s and a DM mass of 1.2 TeV.} \\label{fig:satlocalratio} \\end{figure*}  Many constraints on light-mediator models have been obtained under this assumption that the smooth halo dominates the local DM annihilation signal.   In other words, model parameters for which the smooth halo annihilation signal reproduces the results of PAMELA and \\emph{Fermi} are used to normalize signals in constraining channels where no or little excess is seen.  Accounting for a substructure contribution to the local signal significantly changes the model parameters at which these  constraints should be studied, particularly for mediator masses in the $\\sim 1-200$ MeV range where the local signal is naturally dominated by Sommerfeld-enhanced substructure.  As we shall see, a wide range of constraints become dramatically weaker in this case, specifically: \\begin{itemize} \\item Bounds on the cosmological DM annihilation rate between recombination and reionization, from the anisotropy power spectrum of the cosmic microwave background (CMB), \\item Constraints on light mediators from DM self-interactions, \\item Tensions between the relic abundance and the large local signal, \\item Non-observation of inverse Compton scattering (ICS) and final-state radiation (FSR) signals from annihilation in the Galactic center. \\end{itemize} The parameter space with force carrier masses $\\sim 1-200$ MeV, which is effectively ruled out in the smooth-halo-dominated case, is re-opened in the presence of an $\\mathcal{O}(1)$ substructure contribution to the DM density-squared. This is an important point. Terrestrial fixed-target experiments and low-energy colliders can probe direct production of light hidden-sector gauge bosons \\cite{Pospelov:2008zw}.  Of these, collider searches \\cite{Batell:2009yf,Essig:2009nc,Reece:2009un,:2009pw} can cover the widest range of gauge boson masses, while fixed-target experiments are sensitive to the widest range of couplings, but at lower masses \\cite{Reece:2009un,Bjorken:2009mm,Batell:2009di}. Most fixed-target results \\cite{Bergsma:1985qz,Riordan:1987aw,Bjorken:1988as,Bross:1989mp,Andreas:2010tp,Archilli:2011nh,Merkel:2011ze,  Abrahamyan:2011gv} and recent proposals \\cite{Proposal,Essig:2010xa,HPS,Wojtsekhowski:2009vz,Freytsis:2009bh, HIPS} have the greatest reach at masses below $\\sim 200$ MeV, the region where substructure effects can be most dramatic.  Understanding the role of substructure is essential to clarify what regions of parameter space have astrophysical motivation.  In this note, we will explore the added parameter space that is opened at light ($\\sim$ 1-200 MeV) mediator masses when substructure is taken into account. In Sec. \\ref{sec:constnosub}, we outline the key features of the Sommerfeld enhancement and simple models that illustrate them, and then review previously studied constraints on the relevant parameter space from bounds on DM self-annihilation and self-interaction. We do \\emph{not} initially consider constraint channels that depend on the amount or distribution of substructure, deferring that discussion for later.  Sec. \\ref{sec:effectsofsub} presents a simple parameterization of the boost to the annihilation rate in the presence of both local substructure and Sommerfeld enhancement, and shows how the interplay of these effects relaxes the constraints of Sec. \\ref{sec:constnosub}. We study the parameter space favored to fit the PAMELA and \\emph{Fermi} cosmic-ray signals as a function of the amount of substructure and the range of the DM self-interaction, the effect of the various constraints on this parameter space, and the effect of including substructure on the maximum possible annihilation rate (consistent with all constraints). In Sec. \\ref{sec:constwsub}, we review and address constraints that depend on the distribution of substructure, both in the Milky Way (MW) and on cosmological scales.  These include limits from diffuse $\\gamma$ rays, arising from substructure in the outer MW and in other halos, and limits from the inner MW, which depend upon the extent to which substructure persists in the inner galaxy. ",
        "conclusions": "The $e^+$ excess observed by PAMELA and now \\emph{Fermi} points to a new primary source of cosmic ray electrons and positrons. One of the most exciting, if speculative, explanations of the excess is that it arises from Sommerfeld-enhanced DM annihilation. Such models naturally provide annihilation rates much larger than would be expected from a thermal WIMP with substructure enhancement alone.  Nonetheless, in models with Sommerfeld enhancement, in the presence of $\\mathcal{O}(1)$ substructure, the substructure is often the {\\em dominant} source of the signal, because of the low velocity dispersion of the bound subhalos. Most constraints on these models have been calculated assuming the signal arises from the smooth halo, and the limits become dramatically weaker if substructure dominates the CR signal.  In particular, the constraints on parameter space from the CMB are removed for $\\Delta \\gtrsim 0.4$, since the local signal as well as the early universe signal are both determined by the saturated cross section, in contrast to the smooth halo piece that is generally unsaturated. Because substructure is depleted in the inner regions of galaxies, constraints from FSR and ICS signals in the inner MW are strongly suppressed. As $\\Delta$ increases, lower couplings between the DM and the force carrier are required to fit the CR excesses, and this in turn relaxes limits on the force carrier mass from bounds on DM self-interaction; such constraints are subsumed by the CMB bounds, but depending on the model, the more stringent requirement that self-interaction have negligible effect on dwarf galaxy structure may imply $m_\\phi \\gtrsim 10$ MeV. This has profound implications for terrestrial searches. Including the effects of substructure not only opens the low-mass ($m_\\phi \\sim 1-200 \\mev$) region of parameter space, but possibly makes it the preferred range, once additional constraints are considered. Given the sensitivity that many experiments have in this region, these searches become even more motivated.  Limits on the DM annihilation rate from measurements of the diffuse extragalactic gamma-ray background, or from gamma-rays from the outer halo of the MW, generally become stronger as the amount of substructure is increased. However, the key quantity is the ratio of the  signal in these constraining channels to the local substructure-enhanced annihilation rate, and this can easily be substantially \\emph{smaller} at higher $\\Delta$ than for $\\Delta=0$, relaxing the limits on DM explanations for the CR excesses.  There is no established consensus on what the local boost from substructure should be, due to the large uncertainties in extrapolating the results from $N$-body simulations below their mass resolutions, but it is not thought to be extremely large. Nonetheless, in the presence of Sommerfeld enhancement the substructure contribution can still easily dominate and open up new regions of parameter space, especially for sub-200 MeV force carriers. These regions, too, will be tested by further observations, with the exciting possibility that we might already be detecting the cosmic ray signals of DM substructure."
    },
    "1107/1107.4961_arXiv.txt": {
        "abstract": "In this paper we examine the profiles of the complex Na\\,\\textsc{i}~D-lines in the Red Rectangle. The spectra were acquired with the ARCES \\'{e}chelle spectrograph ($R = 38,000$) on the 3.5-m telescope at the Apache Point Observatory. Additional spectra taken with STIS were acquired from the Hubble Legacy Archive (HLA) and were used to independently confirm the spatial origin of the spectral features. The profile of a single D-line consists of double-peaked emission, red-shifted absorption and blue-shifted absorption. We find that the double-peaked emission originates from the bipolar outflow, the red-shifted absorption feature is due to the photospheric line, and the blue-shifted absorption arises from the bipolar outflow as seen against the photosphere of the luminous post-AGB component in HD~44179. In order to better understand the Na\\,\\textsc{i}~D-line profile, we examined the periodically variable asymmetric photospheric absorption lines.  The asymmetric lines are interpreted as a signature of slow self-accretion following enhanced mass-loss around periastron.  An empirical model was constructed to remove the photospheric component from the Na\\,\\textsc{i}~D-line profile in order to study the nebular emission and absorption of sodium along the line-of-sight to the primary. This paper also discusses the different origins of the single-peaked emission, the double-peaked emission and the blue-shifted and red-shifted absorption components. ",
        "introduction": "The Red Rectangle (RR), an enigmatic protoplanetary nebula, reveals itself to be a complex astrophysical system.  The central source of the nebula, HD~44179, is a binary system in a short-lived evolutionary phase between the asymptotic giant branch (AGB) and the planetary nebula stage of its life, when the expelled circumstellar matter is ionized by the photometric primary star's hot core \\citep{vanwinckel2003, vanwinckel2007, siodmiak2008, szczerba2007}.  The dominant features of the nebula's form are the $\\times$-shaped biconical outflows (nearly 2{\\arcmin} on the sky \\citep{cohen2004}) that emerge from a disc (roughly 0.1{\\arcsec} on the sky), seen nearly edge-on.      We will refer to this optically-thick disc, resolved by \\textit{HST} \\citep{cohen2004} and ground-based interferometric observations \\citep{roddier1995, osterbart1997}, as the circumbinary disc.  The dimensions assumed  for the circumbinary disc are based on a distance of 710~pc \\citep{menshchikov2002}.  While the distance is uncertain, we adopt this value in our investigation.  \\citet{bujarrabal2005} traced the outer radius of the circumbinary disc to 1,850~AU.  The radius of the central cavity in the circumbinary disc is obtained from a scattered light model to be 14~AU \\citep{menshchikov2002}.  The thickness of the circumbinary disc has been determined to be between 90~AU \\citep{menshchikov2002} and 107~AU \\citep{bujarrabal2005, roddier1995, osterbart1997}.  It is worth noting that values of the viewing angle vary from $\\sim$ 0$\\degr$ to 11$\\degr$ \\citep{lopez1995,roddier1995,menshchikov2002}.  These viewing angles would correspond to inclination angles between $\\sim$ 90$\\degr$ to 79$\\degr$.  \\citet{bujarrabal2005} find that an intermediate value of 5$\\degr$ fits their data best, which would correspond to an inclination angle of 85$\\degr$.  At the centre of the cavity in the circumbinary disc are the components of the binary system: the photometric primary (a mass-losing post-AGB star), and the secondary (a main sequence star with an accretion disc and high-velocity jet).  For simplicity we shall refer to the photometric primary as the primary star throughout this paper.  The primary star has an effective temperature of 7,700~K, an estimated mass of 0.8~M$_{\\odot}$, a luminosity of about 6,000~L$_{\\odot}$ \\citep{menshchikov2002}, and a radius of 0.21~AU \\citep{witt2009}.  The secondary has a mass of roughly 0.9~M$_{\\odot}$ and is believed to be a main sequence star \\citep{witt2009}.  In the work of \\citet{witt2009} the dependence of the H$\\alpha$ line on orbital phase was investigated, revealing a high velocity jet ($v_{max} \\sim~560$~km~s$^{-1}$) originating from an accretion disc around the secondary.  The accretion disc around the secondary is powered by mass transfer from the primary star via Roche-lobe overflow at periastron \\citep{witt2009}.  The ionizing photons from the hot inner regions of the accretion disc provide the necessary energy to excite the Extended Red Emission (ERE).  The ERE requires far-ultraviolet photons with energies E $>$ 10.5 eV \\citep{witt2006}.   The accretion disc also provides the energy source to maintain the H\\,\\textsc{ii} region inside the cavity of the optically-thick circumbinary disc that was suggested by \\citet{jura1997}.  The near edge-on optically-thick circumbinary disc prevents a direct line-of-sight view of the binary.  Light originating inside the cavity can only be observed via scattering off material located above and below the disc.  This obscured view of the binary, illustrated in Fig.~\\ref{fig1}, will be referred to as the indirect line-of-sight and is responsible for the difficult interpretation of the spectra.  The indirect line-of-sight to the central object was first studied by \\citet{waelkens1996}, who determined the effective inclination angle to be 35$\\degr$.  The effective inclination angle is the angle at which light from the central binary is scattered into our line-of-sight by the material near the top and bottom of the disc.  This viewing geometry has been confirmed by high-resolution optical \\citep{osterbart1997, cohen2004} and infrared imaging \\citep{tuthill2002}.  For light from the central source (which can only been seen via scattering), we adopt the effective inclination angle of 35$\\degr$ in our calculations.  The outer regions of the nebula are, however, seen directly.  To properly study emission processes that originate in the outer nebula one must use the inclination angle of the nebula (85$\\degr$) rather than the effective inclination angle (35$\\degr$).  Since HD~44179 can only be studied via scattered light, it is important to discuss the known uncertainties. The effective angle at which we are viewing the stars is likely the mean scattering angle. The variation of the scattering angle over the duration of the orbit, which best fits the 0.14~mag variation in the light curve, is 1.6$\\degr$ \\citep{waelkens1996}. The average value for the scattering angle corresponds to the effective inclination angle of 35$\\degr$. In addition, the tilt of the system implies that the scattering angle for light from the north and the south of the nebula is slightly different (of order 5$\\degr$). Therefore, the effective inclination angle is just that, an effective value that we use in order to make progress in understanding this complex system. The variation in the value for the effective inclination angle will cause uncertainties in; the orbital elements, the derived velocities, and the mass of the secondary.  The properties of the stars also depend upon the estimation of the luminosity \\citep{menshchikov2002}. The luminosity depends on the uncertain distance to the star, assumptions about the geometry of the circumbinary-disc, and the evolutionary state of the system. Since this a close binary, it is likely that there has been some interaction in the past resulting in a complex evolutionary history. Therefore, it is likely that the evolutionary track of the stars differ from that of single stars, as in \\citet{Vassiliadis1994} and \\citet{bloecker1995}. As a result, the quoted values that we have used for the mass and luminosity of the primary remain uncertain. At present, these uncertainties can not be addressed with current observational data.   Studying the nebular outflow and its mechanisms is important for understanding how stellar material is returned to the interstellar medium.   The detectable emission of the H$\\alpha$ recombination line in our spectra is limited to the inner parts of the nebula.  The H$\\alpha$ emission is not observed in the bi-conical lobes far away from the central source.  Studying the Na\\,\\textsc{i}~D-lines, easily excited resonance lines, allows us to probe the mass-loss to much greater distances from the central source.  The Na\\,\\textsc{i}~D-lines also allow us to directly study the mass-loss from the primary star.  The details of the observations, and the radial velocity (RV) measurements are given in {\\S}~\\ref{sec-observation}.  The determination of the orbital parameters is discussed in {\\S}~\\ref{sec-orbital}.  In {\\S}~\\ref{sec-abundance} we discuss the Na\\,\\textsc{i} abundance and the model atmosphere used to determine the photospheric line profiles.  The behaviour of the photospheric lines and the processing of light in the RR is discussed in {\\S}~\\ref{sec-lineshape}.  The origins and measurements of the components of the Na\\,\\textsc{i}~D-line profile are discussed in {\\S}~\\ref{sec-naprofiles}.  In {\\S}~\\ref{sec-singlevsdouble} the origin of the single and double peaks are discussed. The conclusions are presented in {\\S}~\\ref{sec-conclusions}.  \\begin{figure}  \\includegraphics[width=80mm]{fig1.eps}  \\caption{This to-scale diagram of the inner region of the RR was constructed for the distance of 710~pc.  The outer radius of the circumbinary disc, 1,850~AU, is not shown.  The central binary is displayed a factor of 10 larger than the scale for the rest of the system.  The inclination angle, angle between the plane of the sky and the equatorial plane, is 85$\\degr$.  The direct and indirect lines-of-sight are also indicated for both the northern and southern ends of the nebula.  The cavity is filled with the low-velocity, mass-loss driven outflow of the primary star and the high-velocity jet from the accretion disc of the secondary.  The location of the primary star is shown for the orbital phase of inferior conjunction; see Fig.~\\ref{fig2} for a to-scale diagram of the binary orbits depicting the orbital phases. \\label{fig1}} \\end{figure} ",
        "conclusions": "\\label{sec-conclusions}  The complex profiles of the Na\\,\\textsc{i}~D-lines in HD~44179, periodically varying in intensity and shape with the orbital period as the photometric primary post-AGB component of the binary moves through its eccentric orbit, reveal detailed information about the varying mass-loss from the primary. Peak mass-loss occurs during the orbital phases around periastron, to be replaced by partial re-accretion during the phases near apastron.  The primary conclusions are:  \\begin{enumerate}  \\item The double-peaked sodium emission lines are produced in the bipolar outflow, seen via the direct line-of-sight (inclination of 85$\\degr$). The observed velocity difference of 12~km~s$^{-1}$ is a result of the 5$\\degr$ tilt of the bipolar outflow axis with respect to the plane of the sky.  \\item The same outflow seen in the double-peaked emission is also seen in absorption against the photospheric continuum of HD~44179 via the indirect line-of-sight (effective inclination angle of 35$\\degr$).  The velocity of the outflow approaching the escape velocity from the primary suggests that the primary star is the ultimate source of the outflowing material.  \\item The single peaked emission lines, for example H$\\alpha$, likely originate from the H\\,\\textsc{ii} region located in the cavity of the circumbinary disc.  The material with single peaked emission lines is at rest with respect to the centre-of-mass.  \\item The red-shifted absorption component of the sodium profile, present for only part of the orbit near periastron, is most likely produced by the underlying photospheric Na\\,\\textsc{i}~D absorption-line of HD~44179.  The specific orbital behaviour is governed by the RV of the primary star.  \\item The photospheric asymmetric blue-winged profiles in HD~44179 are consistent with enhanced mass-loss near periastron passage.  Some of the material lost is accreted onto the disc surrounding the secondary star.  \\item The photospheric asymmetric red-winged profiles are consistent with self-accretion of a portion of the material lost near periastron passage.  This self-accretion occurs on a time scale of half the orbital period, centred around the phase of apastron.  \\end{enumerate}   It is clear that the cause of the broadened photospheric absorption lines is of great importance for understanding and interpreting the spectra of HD~44179.  With current observations, some ambiguities cannot be resolved.  The broadening mechanisms investigated in this paper are either unlikely or cannot fully explain the observed line-profiles.  It is likely that there may be a superposition of mechanisms at play in shaping these spectral features.  The indirect line-of-sight adds to the complication of interpreting the spectra.  The effective inclination angle used here may in fact be somewhat different.  It is more likely that we see HD~44179 under a small range of angles, of which the effective inclination angle is the mean.  It is likely the case that the spread of viewing angles leads to broadened profiles.  In addition, there is a slightly different viewing geometry for the north and south ends of the nebula.  Spatially-resolved, higher-resolution spectra of the photospheric lines could help resolve some of these remaining puzzles.  A full 3-D model is called for in order to understand our complex view of the binary deep inside the circumbinary disc."
    },
    "1107/1107.1069_arXiv.txt": {
        "abstract": "\\noindent This paper presents the analysis of a deep near-infrared $JHK_s$ imaging survey ($37.5\\Box\\dgr$) aimed at tracing the galaxy distribution of the Great Attractor (GA) in the Zone of Avoidance along the so-called Norma Wall. The resulting galaxy catalog is complete to {\\sl extinction-corrected} magnitudes $K_s^o < 14\\fm8$ for extinctions less than $A_K \\le 1\\fm0$ and star densities below $\\log N_{(K<14.0)} \\le 4.72$.  Of the 4360 catalogued galaxies, 99.2\\% lie in the hereby constrained 89.5\\% of the survey area. Although the analyzed galaxy distribution reveals no new major galaxy clusters at the GA distance (albeit some more distant ones), the overall number counts and luminosity density indicate a clear and surprisingly smooth overdensity at the GA distance that extends over the whole surveyed region. A mass estimate of the Norma Wall overdensity derived from (a) galaxy number counts and (b) photometric redshift distribution gives a lower value compared to the original prediction by Lynden-Bell et al. 1988 ($\\sim 14$\\%), but is consistent with more recent independent assessments. ",
        "introduction": "\\noindent The motivation for this deep near-infrared (NIR) imaging survey along the Norma Wall (henceforth the Norma Wall Survey, NWS) is described by Wakamatsu 2011 (these proceedings). It includes a description of the survey area and the derived NIR ($JHK_s$) positional and photometric parameters of the 4360 identified galaxies. The main goal of the project is the mapping of the galaxy distribution in the Great Attractor (GA) region where extinction and star crowding by the Milky Way is so severe that neither optical nor NIR whole-sky surveys (like 2MASX; Jarrett et al. 2000; Jarrett 2004) allow an estimate of the GA mass-density, wheras the deepest systematic HI ZOA surveys to date (Parkes MB; e.g. Kraan-Korteweg et al. \\,2005), which allow full penetration of the ZOA, remain very shallow ($\\la 1$gal$/\\Box\\dgr$). See Kraan-Korteweg \\& Lahav (2000); Kraan-Korteweg (2005) for a review of the various dedicated multi-wavelength ZOA surveys.  The earlier optical, NIR and HI ZOA results suggest the GA to consist of a Great Wall-like structure, that extends over $100\\dgr$ on the sky (from Pavo to Abell S0639); it includes the massive clusters Norma A3627 (the likely core of this structure), CIZA J1324.7-5736 (CIZA), Cen-Crux and PKS 1343-601 (see e.g. Radburn-Smith et al. 2006 for further details).  We used the InfraRed Survey Facility (IRSF) at the 1.4\\,m Japanese telescope in Sutherland (SAAO) to obtain $\\sim 2800$ $JHKs$-band survey images of $7\\farcm8 \\times 7\\farcm8$ to probe the most central, low-latitude ($|b| < 7\\dgr$) region of the Norma Wall. It starts above the Norma cluster (covered by Skelton et al. 2009) and includes the CIZA and Cen-Crux clusters on the other side of the ZOA. The exposure times of 10 min are an optimization between survey depth versus sky background and star density. The longer integrations times and particularly the higher resolution of the IRSF ($45\\arcsec/$pix) compared to 2MASX allows the detection of galaxies deeper into the Milky Way and improved photometry (see also Skelton et al. 2009; Williams et al. these proceedings). Only 5.5\\% of the identified galaxies have counterparts in 2MASX. Further details are given in Riad (2010, PhD thesis) and the forthcoming catalogue paper (Nagayama, Riad et al., in prep). ",
        "conclusions": "The homogeneous galaxy distribution over the survey footprint is suggestive of a continuous structure across the GP, i.e. the so-called Norma Wall. The Norma Wall survey did not uncover any previously unknown clusters that form part of the GA overdensity. Assuming the Great Attractor, respectively the Norma Wall, to be a cylindrical filament of 84.5$h^{-1}$ Mpc in length with a radius of 2$h^{-1}$ Mpc that is filled uniformly by the mean overdensity derived at the approximate GA distance, and including the Norma, Pavo II, CIZA J1324.7-5736, Cen-Crux and Abell S0639 clusters, the total mass excess in galaxies amounts to $\\sim 4\\times 10^{15}h^{-1}\\msun$.  This mass is in good agreement with recent estimates for the mass of the Norma Wall from optical, X-ray and HI observations (Radburn-Smith et al. 2006; Kocevski et al. 2007; Stavely-Smith et al 2000) but lower than the original estimate by Lynden-Bell et al. (1988).\\\\~\\\\   \\noindent {\\bf Acknowledgements. ---} \\begin{small} {RCKK and IFR acknowledge financial support from the National Research Foundation through a mobility grant from the South African - Japanese (NRF/JSPS) bilateral grant in Astronomy. IFR was supported throughout his PhD studies by the University of Khartoum, South African National and the Stichting Steunfonds Soedanese Studenten. We are grateful to T. Jarrett for his derivation of the photometric redshifts.} \\end{small}"
    },
    "1107/1107.3770_arXiv.txt": {
        "abstract": "The X-ray source 1E\\,161348--5055 lies at the centre of the 2-kyr-old supernova remnant \\snr. Owing to its 24-ks modulation, orders-of-magnitude flux variability over a few months/years, and lack of an obvious optical counterpart, 1E\\,161348--5055 defies assignment to any known class of X-ray sources. Starting from April 2006, \\swift\\ observed 1E\\,161348--5055 with its X-ray telescope for $\\sim$2 ks approximately once per month. During the five years covered, the source has remained in a quiescent state, with an average observed flux of $\\sim$$1.7\\times10^{-12}$ \\flux\\ (1--10 keV), $\\sim$20 times lower than the historical maximum attained in its 1999--2000 outburst. The long time-span of the \\swift\\ data allows us to obtain an accurate measure of the period of 1E\\,161348--5055 [$P=24\\,030.42(2)$ s] and to derive the first upper limit on its period derivative ($|\\dot{P}|<1.6\\times10^{-9}$ s s$^{-1}$ at 3$\\sigma$). ",
        "introduction": "1E\\,161348--5055 (hereafter \\src) was discovered with the \\emph{Einstein} satellite \\citep{tuohy80} close to the geometrical centre of the young supernova remnant (SNR) \\snr\\ (age $\\sim$2 kyr; \\citealt*{carter97}). It was proposed as the first example of a radio-quiet (possibly owing to an unfavourable radio beaming), isolated, cooling neutron star (\\citealt{tuohy80,tuohy83}; \\citealt*{gph97}).  At present, there is little doubt that \\src\\ indeed is a neutron star \\citep{deluca06} and the source is traditionally included in the class of the `central compact objects' (CCOs; see \\citealt{deluca08} for a review). CCOs are a small group of young and seemingly isolated X-ray-emitting neutron stars (with thermal-like spectra), observed close to the centre of non-plerionic SNRs and without obvious counterparts in other wavebands. However, its peculiar temporal behaviour distinguishes \\src\\ from the other CCOs (actually, it singles this source out as a unique object in general). The first peculiarity of the source is its orders-of-magnitude X-ray flux variability on a few months/years time-scale (\\citealt*{gotthelf99}; \\citealt{garmire00,sanwal02,becker02}). Moreover, the first \\cxo\\ observation of \\src\\ in a low state hinted at a possible periodicity at $\\sim$6 hours \\citep{gpg00} that was not confirmed by subsequent observations of the source in bright states. A long (90 ks) observation with \\xmm, performed in 2005, caught \\src\\ in a low state and yielded unambiguous evidence for a strong, nearly sinusoidal modulation at $6.67\\pm0.03$ hours ($24.0\\pm0.1$ ks; \\citealt{deluca06}). The same periodicity was then recognised also in the older data-sets, albeit with a very different pulse shape, including two narrow dips per period. No faster pulsations are seen in \\src\\ \\citep{deluca06}.  Large flux variations, similar to those observed in \\src, are common among magnetars (e.g. \\citealt{rea11}), but these pulsars, whose emission is believed to be powered mainly by the magnetic field, are characterised by rotational periods in the narrow range 2--12 s. On the other hand, CCOs are steady sources, and their periods --when known-- are in the 0.1--0.5 s range (\\citealt{zavlin00}; \\citealt*{ghs05,gotthelf09}). If \\src\\ is indeed a magnetar, it must have been slowed down by some unusual mechanisms, perhaps by a propeller interaction with a debris disk \\citep{deluca06,li07}. A different possibility is that \\src\\ is a peculiar low-mass binary,\\footnote{Deep observations of the field of \\src\\ with the Very Large Telescope and the \\emph{Hubble Space Telescope} showed only two or three faint infrared sources ($H\\sim22$) consistent with the position of \\src. If none of them is linked to the X-ray source, \\src\\ is undetected in the near infrared down to $H>23$ \\citep{dmz08}.} powered by a double (wind plus disk) accretion onto a recently-born compact object (in this case the 24-ks signal would result from the orbital motion of the system) or hosting a magnetar \\citep{deluca06,pizzolato08,bhadkamkar09}. Both scenarios require nonstandard assumptions about the formation and evolution of compact objects in supernova explosions.  Here we report on the results from a 5-year monitoring of \\src\\ with the \\swift\\ satellite \\citep{gehrels04short}. This unique data set allowed us to obtain a phase-coherent timing solution encompassing also \\cxo\\ and \\xmm\\ archival observations. Thanks to this, we are able to derive an accurate period for \\src\\ and to set the first upper limits on the period derivative for this puzzling source. ",
        "conclusions": "We presented the analysis of the first five years (2006 April--2011 April) of the \\swift/XRT monitoring of the enigmatic X-ray source \\src\\ at the centre of the SNR \\snr. During this time span, the source remained in a quiescent state, with an average observed 1--10 keV flux of $\\sim$$1.7\\times10^{-12}$ \\flux, 20--30 times lower than the historical maximum attained in the 1999--2000 outburst ($\\sim$$5\\times10^{-11}$ \\flux; \\citealt{garmire00}). The timing study of the \\swift\\ data yielded an accurate measure of the modulation period of \\src\\ which allowed us to phase-connect the XRT data with archival \\cxo\\ and \\xmm\\ observations. We derived two timing solutions, labelled A and B, both consistent with the same constant period $P_{\\mathrm{A}}=P_{\\mathrm{B}}=24\\,030.42(2)$ s but with somewhat different upper limits on the period derivative. Solution A ($|\\dot{P}_{\\mathrm{A}}|<3.3\\times10^{-9}$ s s$^{-1}$; MJD 53\\,605--55\\,632) is based on the single-peaked pulse profiles observed while \\src\\ was in quiescence. Solution B ($|\\dot{P}_{\\mathrm{B}}|<1.6\\times10^{-9}$ s s$^{-1}$; MJD 52\\,336--55\\,632) is a natural extension of solution A including the multi-peaked \\cxo/ACIS-S profile obtained on 2002 March 03, when \\src\\ was in a rather bright state \\citep{sanwal02}.  The phase-coherent timing technique employed in our timing study of \\src\\ closely parallels the well-tested (virtually all systematics are under control) procedures used for X-ray bright spin-powered pulsars. For such objects it is implicitly assumed that the pulse profile does not intrinsically change in time, so that, when a Fourier decomposition of the pulse is introduced, neither the fundamental nor the higher harmonics evolve. In such a picture all the variability is due to Poisson noise and the phase evolution can be tracked by matching the pulse profile at any given epoch always with the same template.  In the case of \\src\\ there are clear indications that the pulse shape changes in time so that the previous assumptions are not valid. This forced us to abandon the standard template-matching analysis and follow instead the phase of the fundamental (first) harmonic.\\footnote{The large phase uncertainties do not allow a similar study of the higher harmonics (which are not even always detectable).} The main justification for such an approach is that usually (moderate) pulse shape changes are associated with higher harmonics while the fundamental is stable. Even if this seems to be the case for \\src, given the stability of the phase of the fundamental harmonic over several years, we stress that it does not have to hold in general (see, e.g., \\citealt{hartman08}).  In the following we discuss the implications of our newly derived upper limit $\\vert\\dot P\\vert < 1.6\\times 10^{-9}$ s s$^{-1}$ on the models that have been proposed so far for \\src. We remark that all the ensuing considerations are based on the assumption (discussed above) that the fundamental harmonic is a good tracer of the timing behaviour of the source.  Although the nature of \\src\\ is still an open issue, most interpretations favour the neutron star scenario, in which the star is either isolated \\cite[][]{deluca06,li07} or in a binary system with a low-mass companion \\cite[][]{pizzolato08,bhadkamkar09}. While the model by \\cite{bhadkamkar09} is based on a fast-spinning, moderately magnetised neutron star in a 6.67 hr eccentric orbit, in all the other cases an ultra-magnetised neutrons star ($B\\approx 10^{15}\\ {\\rm G}$) is required to explain the very long period of \\src\\ and the observed X-ray periodicity is related to the star spin period (in the binary scenario of \\citealt{pizzolato08} the latter may or may not coincide with the orbital period). Actually, as already noted by \\cite{deluca06}, magneto-dipolar braking alone is not enough to explain the present value of $P$ even invoking a magnetar, and spin-down by the interaction of the star magnetosphere with a (residual) disc is also necessary for an isolated object. Spin-down to $P\\simeq 6.67\\ {\\rm hr}$ in a time $\\sim$2 kyr can be achieved if the initial period is peculiarly long ($\\ga$300 ms, \\citealt{deluca06}; magnetars are in fact believed to be born with ms periods, \\citealt{thompson93}). Using different assumptions on the termination radius of the disc \\cite[see][]{eksialpar05}, \\cite{li07} showed that the same result can be recovered also for more conventional values of the initial period ($\\approx$10 ms).  The current spin-down rate expected in the binary magnetar model is very small since the equilibrium period is reached well in advance of $\\sim$1000  yr, unless the synchronisation time is very short ($\\sim$10 yr) and there is no mass transfer in the system \\cite[][]{pizzolato08}. If \\src\\ is an object of this kind our upper limit on $\\dot P$ is completely non-constraining for the model.  On the other hand, if \\src\\ is an isolated magnetar surrounded by a fossil disc it must be a quite rare system. According to the Monte Carlo simulations of \\cite{li07}, the fraction of objects of this type with periods $\\ga$100 s at an age of $2.5$ kyr is only $\\sim$1\\%. The large majority of stars are much faster rotators ($P\\sim\\ 1$--10 s) in the ejector stage (typically identified with SGRs/AXPs) while ultra-slow systems are in the propeller/accretor phase.\\footnote{Although previous figures are model-dependent the conclusion that propellers/accretors are a tiny minority appears to be robust.}  The spin-down rate of a neutron star which interacts with a fossil disc in the propeller stage is given by $\\dot P \\sim -\\dot MR_{in}^2[\\Omega_K(R_{in})-2\\pi/P]P^2/(\\pi I)$ where $R_{in}=0.5[B^4R^{12}/(8GM\\dot M^2)]^{1/7}$ is the Alfv\\'en radius, $M\\sim 1.4M_\\odot$ and $R\\sim 10^6\\ {\\rm cm}$ are the star mass and radius ($I\\sim 10^{45}\\ {\\rm g\\,cm}^2$ is the moment of inertia), $B$ is the surface magnetic field, $\\Omega_K$ is the Keplerian angular velocity and $\\dot M\\propto t^{-1.25}$ is the mass loss rate from the disc \\cite[][and references therein]{li07}.\\footnote{We remark that no complete theory of the interaction of a disc with a magnetised neutron star exists. Following \\cite{li07} we adopted the `efficient' form of the propeller torque \\cite[see e.g.][for a comparison of different expressions of the torques]{franci02}.} If \\src\\ is currently in the propeller phase, $\\dot P$ can be derived from the previous expressions with $P\\sim 24$ ks and $t\\sim 2.5$ kyr, as a function of the disc initial mass $M_D(0)$ and of $B$. Results are shown in Fig.~\\ref{pdot} (lower panel) for $10^{-6} < M_D(0)/M_\\odot < 10^{-2}$ and $3.2\\times 10^{13}\\ {\\rm G} < B < 10^{16}\\ {\\rm G}$; the Alfv\\'en radius as a function of $M_D(0)$ and $B$ is plotted in the upper panel. Comparison of $R_{in}$ with the light cylinder radius, $R_{LC}=cP/(2\\pi)$, and the co-rotation radius, $R_{C}=[GMP^2/(4\\pi^2)]^{1/3}$ shows that the star is currently in the propeller phase, characterised by $R_C<R_{in}<R_{LC}$, only if $B\\ga 10^{15}\\ \\rm G$. Lower values of the field or large enough (depending on $B$) initial disc mass would result in \\src\\ being in the accretor stage. The resulting values of $\\dot P$ are orders of magnitude above the upper limit we reported (as a reference, the case discussed in \\citealt{deluca06} has $\\dot P\\sim 10^{-7}$ s s$^{-1}$), with the only exception of an extremely narrow range of $M_D(0)$ for each $B$. Although not all the values of $M_D(0)$ and $B$ we considered will produce the correct period at an age of 2.5 kyr (this depends also on both the initial period and the angle between the magnetic and spin axes, see again \\citealt{li07}), our conclusion is that if \\src\\ is an isolated magnetar, and hence it must be either a propeller or an accretor at present, the typical large expected values of $\\dot P\\gg 3\\times 10^{-9}$ s s $^{-1}$ makes the propeller option rather unlikely. Assessing the reliability of the accretor scenario requires a more thorough analysis. Here we just note that if the same expression of the torque remains valid during the accretion phase, the spin-up rate is also largely in excess of our upper limit on $\\dot P$.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{pdot.eps}} \\caption{\\label{pdot} Upper panel: the Alfv\\'en radius as a function of the initial disc mass $M_D(0)$ for different values of the magnetic field in the range $13.5\\leq \\log B\\, (\\rm G)\\le 16$ (the field increases as the curve moves upwards); the light cylinder radius $R_{LC}$ and the corotation radius $R_C$ are also shown. Lower panel: the period derivative as a function of $M_D(0)$ for those values of $B$ for which the source is currently in the propeller stage; the dashed line marks our derived upper limit on $\\dot P$. Here the period and age have been fixed to those inferred for \\src\\ (see text).} \\end{figure}"
    },
    "1107/1107.4366_arXiv.txt": {
        "abstract": "\\vspace{1pt}  The outer haloes of the Milky Way (MW) and Andromeda (M31) galaxies contain as much important information on their assembly and formation history as the properties of the discs resident in their centres.  Whereas the structure of dark matter (DM) haloes has been studied for a long time, new observations of faint structures hiding in the depths of the stellar halo have opened up the question of how the stellar halo is related to the DM underlying it. In this paper we have used the Constrained Local UniversE Simulation (CLUES) project to disentangle the stellar and DM component of three galaxies that resemble the MW, M31 and M33 using both DM only simulations and DM + gas-dynamical ones. We find that stars accreted in substructures and then stripped follow a completely different radial distribution than the stripped DM: the stellar halo is much more centrally concentrated than DM. In order to understand how the same physical process - tidal stripping  -  can lead to different $z=0$ radial profiles, we examined the potential at accretion of each stripped particle. We found that star particles sit at systematically higher potentials than DM, making them harder to strip. We then searched for a threshold in the potential of accreted particles $\\phi_{\\rm th}$, above which DM particles in a DM only simulation behave as star particles in the gas-dynamical one. We found that in order to reproduce the radial distribution of star particles, one must choose DM particles whose potential at accretion is $\\gsim16\\phi_{\\rm subhalo}$, where $\\phi_{\\rm subhalo}$ is the potential at a subhaloes edge at the time of accretion.  A rule as simple as selecting particles according to their potential at accretion is able to reproduce the effect that the complicated physics of star formation has on the stellar distribution. This result is universal for the three haloes studied here and reproduces the stellar halo to an accuracy of within $\\sim2\\%$. Studies which make use of DM particles as a proxy for stars will undoubtedly miscalculate their proper radial distribution and structure unless particles are selected according to their potential at accretion. Furthermore, we have examined the time it takes to strip a given star or DM particle after accretion. We find that, owing to their higher binding energies, stars take longer  to be stripped than DM. The stripped DM halo is thus considerably older than the stripped stellar halo. ",
        "introduction": "\\label{introduction} According to the current accepted cosmological model, the universe is composed of 26\\% cold dark matter (DM), 70\\% dark energy ($\\Lambda$) and 4\\% baryons \\citep[e.g.][]{2007ApJS..170..377S}. Structure in the so-called $\\Lambda$CDM cosmology forms from the bottom up \\citep{1978MNRAS.183..341W,1985ApJ...292..371D} - the first objects to collapse at high redshift are small sub-galactic units, that condense out of small perturbations in the initial density field. These clumps then merge in an hierarchical fashion to construct the large bound objects we observe in the local universe. Visible luminous matter -  stars - are believed to be formed when Giant Molecular Clouds collapse in the potential wells of these bound blobs of DM and gas.  The merging process gives rise to DM haloes, which today host bright central galaxies such as the Milky Way (MW) and the Andromeda galaxy (M31) in their cores. The outskirts of such DM haloes are populated by a two component medium: diffuse matter and matter bound to substructures. Much of the mass is found bound to satellite galaxies which orbit within their parent halo. The properties (age, orbital parameters, spatial distribution, kinematics, etc)  of luminous satellite galaxies can teach us a lot regarding the formation of their hosts and have been the target of numerous observational and theoretical studies. Indeed, the past 5 years has seen an increased focus on the detection of satellite galaxies and has resulted in around a dozen new satellites being detected by the SDSS \\citep{2008ApJ...686..279K,2009MNRAS.397.1748B}.  Yet the $z=0$ satellite galaxy population is not a full survey of all the substructures accreted by the parent DM halo, since many substructures accreted at high redshift will, by $z=0$, have been tidally disrupted by the host potential, resulting in the stripping of dark matters and stars. Indeed it is believed that the stellar halo - stars exterior to the central galaxy and not bound to substructures - was formed by the tearing of stars from accreted satellite galaxies. \\cite{2010MNRAS.406..744C} argue that the vast majority of stars in the MW's halo were stripped from just one or two large satellites. \\cite{2009ApJ...702.1058Z} have studied the stellar halo in gas-dynamical/N-body simulations and have identified that in fact the stellar halo has a dual-origin: part of it was created via tidal stripping of stars from disrupted satellites, and part was pushed out of central galaxies during minor mergers. In a follow up paper,  \\cite{2010ApJ...721..738Z} argue that the metallic abundance patterns (of [Fe/H] and [O/Fe]) of stars can be used to distinguish between theses different formation mechanisms.  Observations using the SDSS by, e.g., \\cite{2008ApJ...680..295B} have indicated that the MW's halo is consistent with being formed entirely out of accreted debris material. This is in disagreement with observations by \\cite{2008Natur.451..216C} who find clear differences between kinematical properties of the inner and outer stellar halo - stars in the inner halo are found to exhibit a net prograde rotation while the outer halo is dominated by retrograde motion. It has been suggested that the differences in net rotation of stars in the halo betray a dual origin, a result recently supported by \\cite{2011arXiv1104.2513B} who find differences in kinematics for inner and outer halo stars.  The situation with M31 is similar as observations seem to favor an accreted origin. By studying the ages of stars in the halo, \\cite{2008ApJ...685L.121B} argue against a in-situ origin and, because the stars are by and large relatively old - they argue for an hierarchical build up of M31. By focusing on the metal enrichment, \\cite{2009ApJ...701..776G} too seem to argue for an entirely accreted stellar halo with little evidence of in-situ star formation.  The assembly history of the DM halo on the other hand, is more difficult to pin down as direct observations are by definition impossible. Yet many authors \\citep[e.g.][and references therein]{2001ApJ...554..903K,2002ApJ...568...52W,2003MNRAS.339...12Z,2007ApJ...667..859D} have used $N$-body simulations to determine the relative importance of the two main ``modes'' of halo growth: diffuse accretion versus mergers. Most recently  \\cite{2010arXiv1008.5114W} have used the Aquarius simulation and found that ambient accretion contributes the largest amount of material to the dark halo.  It is thus unclear if the conclusions drawn from observations of our stellar halo - that the stars were stripped from infalling satellite galaxies - are consistent with DM only simulations which point towards a diffuse smoothly accreted halo where mergers and debris material play a minor role. In this paper we disentangle these two components in order to understand how they co-evolved. ",
        "conclusions": "\\label{sec:conclusion} In this paper we have studied three Galaxy sized haloes formed in a constrained simulation of the Local Group. The three galaxies have approximately the same size and relative positions (as well as some other properties) as the observed Local Group members and are therefore refereed to as MW, M31 and M33. By construction, the galaxies are formed in an environment whose bulk properties (e.g. distance to a Virgo mass cluster) closely match observations. We use our constrained Local Group to focus on the similarities and differences in the origin of the DM and stellar halo.   We have focused our analysis on attempting to understand why stripped DM particles orbiting in the parent halo, have a completely different radial distribution than ex-situ star particles, who were unbound by the exact same process. Specifically, the stars are more centrally concentrated then the DM, which roughly follows an NFW profile \\citep[e.g.][]{2010arXiv1008.5114W}. Their different $z=0$ distribution implies that the process of tidal stripping must be operating on different subsets of the particles' distribution.  We searched for a subset of the DM particles that is indistinguishable from the stellar halo, in order to understand how these very similar components evolved differently from each other. We found that DM particles (in a gas-dynamical simulation) that had potentials at accretion of at least $\\sim12$ times the potential at the subhalo's edge, are able to closely match the radial distribution of star particles. In a DM only simulation, the threshold above which selected DM particles reproduced the stellar halo distribution was $\\sim 16$ times the potential at the subhaloes edge.   Our threshold value is constant across the three haloes we have studied with very small scatter, despite the fact that our small sample size of three (roughly) Milky Way sized haloes have very different mass accretion histories.  We have thus found a method through which the stellar halo may be modelled without the need of running a semi-analytical model to treat the baryons. By simply selecting those particles at accretion whose potential is greater than the threshold value quoted here, a set of particles can be identified that nearly perfectly matches the stellar distribution.   The implications of this work are therefore that the stripped stellar halo reflects the fate of material that sits deep in a halo's potential well in an absolute sense, not in a relative one. It is important to note that selecting a fraction of DM particles according to their \\textit{relative} binding energy at accretion (e.g. the 10\\% most bound particles) will not successfully reproduce the $z=0$ stellar distribution  unless the 10\\% most bound particles sit within a region of the potential which is greater than $\\sim 16$ the potential at the host's edge. In fact selecting the 10\\% most bound DM particles at accretion returns a halo profile which follows the DM and is around 20\\% less centrally concentrated than the stars. This is because star formation occurs according to a local density criteria, not according to a ``global'' property, such as a particles binding energy relative to the entire subhalo. A given halo will only form stars in its centre if (a) it is large enough to retain its gas and shield it from photo-ionization \\citep[e.g][]{2003ApJ...599...38B} and (b) the potential in its centre is deep enough such that the gas density may trigger star formation. Thus the likelihood of star formation depends just on structural parameters (like concentration) and whether densities in a halo's centre are high enough.   Additionally since the particles that sit in the deepest regions of the subhalo are harder to unbind, they will become unbound later - this implies that the stellar halo is younger then the DM halo: the star particles that compose the stellar halo were unbound from their substructures and deposited in the halo at lower $z$, than the DM. When we observe stars in the outer halo of the MW or M31, we must take care drawing conclusions regarding the DM haloes assembly.   The diffuse DM halo has profoundly different properties to the diffuse stellar halo. Its lack of central concentration dominates its global profile. Although there are some detailed differences among our three haloes in the relative contribution from in-situ and stripped material to the stellar background, the bottom line is that DM particles can not serve as a proxy for star particles unless care is taken in their selection.  Two main results have been presented here. The first is that the stars that constitute the diffuse stellar halos form at the bottom of the potential well of subhaloes, and hence are more bound at infall than the corresponding DM particles that make the diffuse DM halo. This is a 'trivial' fact reproduced by all simulations of galaxy formation, yet there is no general consensus on the details underlying this fact. The novel and non-trivial other results is that one can find a simple mapping that enables the association of a subset of the DM particles with the stellar halo particles. The mapping is based on one single parameter, namely the scaled value of the depth of the potential well within which the halo stars formed. Such a simple mapping of the stellar halo and the subset of DM particles, is a robust outcome of the hierarchical nature of galaxy formation. Yet, the actual value of the parameter that controls the mapping most probably varies with the particular implementation of numerical simulations of galaxy formation. This parameter might vary with the details of how star formation and feedback processes are modeled as well as on the particular numerical schemes applied. This is posed here as an open question that we hope will be addressed by practitioners in the field."
    },
    "1107/1107.4372_arXiv.txt": {
        "abstract": "We jointly constrain the luminosity function (LF) and black hole mass function (BHMF) of broad-line quasars with forward Bayesian modeling in the quasar mass-luminosity plane, based on a homogeneous sample of $\\sim 58,000$ SDSS DR7 quasars at $z\\sim 0.3-5$. We take into account the selection effect of the sample flux limit; more importantly, we deal with the statistical scatter between true BH masses and FWHM-based single-epoch virial mass estimates, as well as potential luminosity-dependent biases of these mass estimates. The LF is tightly constrained in the regime sampled by SDSS, and makes reasonable predictions when extrapolated to $\\sim 3$ magnitudes fainter. Downsizing is seen in the model LF. On the other hand, we find it difficult to constrain the BHMF to within a factor of a few at $z\\ga 0.7$ (with \\MgII\\ and \\CIV-based virial BH masses). This is mainly driven by the unknown luminosity-dependent bias of these mass estimators and its degeneracy with other model parameters, and secondly driven by the fact that SDSS quasars only sample the tip of the active BH population at high redshift. Nevertheless, the most likely models favor a positive luminosity-dependent bias for \\MgII\\ and possibly for \\CIV, such that at fixed true BH mass, objects with higher-than-average luminosities have over-estimated FWHM-based virial masses. There is tentative evidence that downsizing also manifests itself in the active BHMF, and the BH mass density in broad-line quasars contributes an insignificant amount to the total BH mass density at all times. Within our model uncertainties, we do not find a strong BH mass dependence of the mean Eddington ratio; but there is evidence that the mean Eddington ratio (at fixed BH mass) increases with redshift. ",
        "introduction": "\\label{sec:intro}  One major effort in modern galaxy formation studies is to understand the cosmic evolution of supermassive black holes (SMBHs), given their ubiquitous existence in almost every local bulge-dominant galaxy, and possible roles during their co-evolution with the host galaxy \\citep[e.g.,][]{Magorrian_etal_1998,Gebhardt_etal_2000a,Ferrarese_Merritt_2000,Gultekin_etal_2009, Hopkins_etal_2008a,Somerville_etal_2008}. Over the past decade, the rapidly growing body of observational data and numerical simulations have led to a coherent picture of the cosmic formation and evolution of SMBHs within the hierarchical $\\Lambda$CDM paradigm \\citep[e.g.,][]{Haiman_Loeb_1998,Kauffmann_Haehnelt_2000,Wyithe_Loeb_2003,Volonteri_etal_2003,Hopkins_etal_2006, Hopkins_etal_2008a,Shankar_etal_2009a,Shankar_etal_2010,Shen_2009}. Although many fundamental issues regarding SMBH growth still remain unclear (such as BH seeds, fueling and feedback mechanisms), these cosmological SMBH models are starting to reproduce a variety of observed SMBH statistics in an unprecedented manner.  It is now widely appreciated that SMBHs grow by gas accretion in the past, during which they are witnessed as quasars and active galactic nuclei (AGNs) \\citep[e.g.,][]{Salpeter_1964,Zeldovich_Novikov_1964,Lynden-Bell_1969}. In the local Universe, the mass function of dormant SMBHs is estimated by convolving the galaxy distribution functions with various scaling relations between galaxy properties and BH mass. This relic SMBH population has been used to constrain the accretion history of their active counterparts, using the So{\\l}tan argument and its extensions \\citep[e.g.,][]{Soltan_1982,Small_Blandford_1992,Salucci_etal_1999,Yu_Tremaine_2002,Yu_Lu_2004,Yu_Lu_2008,Shankar_etal_2004,Shankar_etal_2009a, Marconi_etal_2004,Merloni_2004,Hopkins_etal_2007a,Merloni_Heinz_2008}. The agreement between the relic BH mass density and the accreted mass density provides compelling evidence that these two populations are ultimately connected. Therefore it is of imminent importance to quantify the abundance of active SMBHs as a function of redshift.  The demographics of the active SMBH population has been the central topic for quasar studies since the first discovery of quasars \\citep[][]{Schmidt_1963,Schmidt_1968}. Traditionally this is done in terms of the luminosity function (LF), i.e., the abundance of objects at different luminosities. Measuring the LF and its evolution has been the most important goal for modern quasar surveys \\citep[e.g.,][]{Schmidt_Green_1983,Green_etal_1986}. In the last decade the LF has been measured for different populations of active SMBHs and in different bands \\citep[e.g.,][]{Fan_etal_2001,Fan_etal_2004,Boyle_etal_2000, Wolf_etal_2003,Croom_etal_2004,Croom_etal_2009,Hao_etal_2005b, Richards_etal_2005,Richards_etal_2006a,Jiang_etal_2006,Jiang_etal_2008,Jiang_etal_2009, Fontanot_etal_2007,Bongiorno_etal_2007, Willott_etal_2010b, Ueda_etal_2003,Hasinger_etal_2005,Silverman_etal_2005,Silverman_etal_2008,Barger_etal_2005}, and it constitutes a crucial observational component in all cosmological SMBH models.  A more important physical quantity of SMBHs is BH mass. BH mass is directly related to growth, and when the BH is active, it determines the accretion efficiency ($\\dot{M}/M_{\\rm BH}$) via the Eddington ratio and an assumed radiative efficiency (e.g., such as the average value constrained by the Soltan argument). Thus knowing the mass function of SMBHs as a function of redshift adds significantly to our understandings of their cosmic evolution.  It remains challenging to directly measure the dormant BHMF at high redshift. This is not only because the galaxy distribution functions are less well-constrained at high redshift, but also because the evolution of the scaling relations (both the mean relation and the scatter) between galaxy properties and BH mass is poorly understood. On the other hand, it has become possible to measure the active BHMF of broad-line quasars\\footnote{From now on, unless otherwise specified, we use the term ``quasar'' to refer to broad-line (type 1) quasars for simplicity.}, using the so-called virial BH mass estimators based on their broad emission line and continuum properties measured from single-epoch spectra \\citep[e.g.,][]{Wandel_etal_1999,Mclure_Dunlop_2004,Vestergaard_Peterson_2006}, a technique rooted on reverberation mapping (RM) studies of local broad-line AGNs \\citep[e.g.,][]{Blandford_McKee_1982,Peterson_1993,Kaspi_etal_2000,Peterson_etal_2004, Bentz_etal_2006,Bentz_etal_2009}. These single-epoch virial BH mass estimators are calibrated empirically using the RM AGN sample to yield on average consistent BH mass estimates compared with RM masses, which are further tied to the BH masses predicted using the $M_{\\rm BH}-\\sigma$ relation \\citep[e.g.,][]{Tremaine_etal_2002,Onken_etal_2004}. The nominal scatter between these single-epoch virial estimates and the RM masses is on the order of $\\sim 0.4$ dex \\citep[e.g.,][]{Mclure_Jarvis_2002,Mclure_Dunlop_2004,Vestergaard_Peterson_2006}.  A couple of recent studies have applied this technique to measure the active BHMF with statistical quasar and AGN samples \\citep[e.g.][]{Greene_Ho_2007a,Vestergaard_etal_2008,Vestergaard_Osmer_2009, Schulze_Wisotzki_2010}. A robust determination of the active BHMF constitutes an important building block of cosmological SMBH models, in addition to the luminosity function. These virial mass estimators also enable statistical studies on the Eddington ratios of broad-line quasars and AGNs \\citep[e.g.,][]{Vestergaard_2004,Mclure_Dunlop_2004,Kollmeier_etal_2006,Sulentic_etal_2006,Babic_etal_2007, Jiang_etal_2007,Kurk_etal_2007,Netzer_etal_2007,Shen_etal_2008b,Gavignaud_etal_2008, Labita_etal_2009a,Trump_etal_2009,Trump_etal_2011,Willott_etal_2010,Trakhtenbrot_etal_2011}, over a wide range of luminosities and redshifts, and therefore provide constraints on the accretion efficiency of these active SMBHs.  With the development of these virial mass estimators, we now have both BH mass estimates and luminosities for broad-line quasar samples. Given the intimate relation between BH mass and luminosity, it is important and necessary to study their joint distribution and evolution in the mass-luminosity plane \\citep[e.g.,][]{Steinhardt_Elvis_2010a,Steinhardt_etal_2011a}. This represents a significant step forward to study the demography of quasars than using LF alone, and offers new insights on the properties and evolution of the active SMBH population.  However, the importance of distinguishing between virial mass estimates and true BH masses can hardly be over-stressed. While these virial estimators currently are the only practical way to estimate BH masses for large samples of broad-line quasars and AGNs, the nontrivial uncertainty of these imperfect estimators has severe impact on the mass distribution under study. The difference between virial masses and true masses not only modifies the underlying true distribution, but also introduces Malmquist-type biases \\citep[e.g.,][]{Shen_etal_2008b,Kelly_etal_2009a,Shen_Kelly_2010}. These effects tend to dilute any potential mass-dependent trends or correlations \\citep[e.g.,][]{Kelly_Bechtold_2007,Shen_etal_2009}, and may lead to unreliable conclusions. Thus it is important to consider these effects when the statistics is becoming good enough.  \\citet{Kelly_etal_2009a} developed a Bayesian framework to estimate the BHMF/LF for broad-line quasars, which accounts for the uncertainty in virial BH mass estimates, as well as the selection incompleteness in BH mass (since the sample is selected in luminosity). This method was subsequently applied to the SDSS DR3 quasar sample \\citep{Kelly_etal_2010}, based on virial mass estimates from \\citet{Vestergaard_etal_2008}. This Bayesian framework is a more rigorous and quantitative treatment than the simple forward modeling performed in \\citet{Shen_etal_2008b}, and allows a more reliable measurement of the true active BHMF and its uncertainty for quasars.  Equipped with an improved version of this Bayesian framework, in this paper we measure the active BHMF and LF based on a homogeneous sample of $\\sim 58,000$ quasars from SDSS DR7 with FWHM-based virial mass estimates from \\citet{Shen_etal_2011a}. The much improved statistics now allows a detailed examination of the joint distribution in the mass-luminosity plane, and provides better constraints on BH accretion properties.  A key difference in our approach compared with most earlier work is the attempt to account for the uncertainty (error) in these virial mass estimates. We distinguish three types of errors in single-epoch virial BH mass estimates: \\begin{itemize}  \\item  measurement error, which is propagated from the uncertainties of FWHM and continuum luminosity measurements from the spectra; the measurement errors are typically $\\ll 0.3$ dex for our sample (see Fig.\\ \\ref{fig:bh_dist}) and hence are negligible; however, measurement error may become important for other samples with low spectral quality.  \\item statistical error, which is the scatter of virial BH masses around RM masses when these virial estimators were calibrated against local RM AGN sample; the statistical error is $\\ga 0.3$ dex \\citep[e.g.,][]{Mclure_Jarvis_2002, Mclure_Dunlop_2004,Vestergaard_Peterson_2006}, which will be taken into account in our Bayesian approach.  \\item  systematic biases, which may result from the virial assumption, the usage of RM masses as true masses during calibration, the usage of a particular definition of line width as the surrogate for the virial velocity, the extrapolation of the virial calibrations to high luminosity/redhift, as well as other possible systematics \\citep[e.g.,][]{Krolik_2001,Collin_etal_2006,Shen_etal_2008b,Marconi_etal_2008, Fine_etal_2008, Fine_etal_2010, Netzer_2009, Denney_etal_2009,Wang_etal_2009b,Graham_etal_2011, Rafiee_Hall_2011b,Steinhardt_2011}.  We generally neglect systematic biases in the current study, as they are poorly understood at present. That means we assume {\\em on average} these virial mass estimators give unbiased mass estimates (see \\S\\ref{subsec:prelim} for the meaning of ``unbiased''). However, we do consider a possible luminosity-dependent bias \\citep[e.g.,][]{Shen_etal_2008b,Shen_Kelly_2010}, which we describe in detail in \\S\\ref{subsec:prelim}. This is not only because luminosity is an explicit term in all virial estimators, but also because that many studies with virial BH masses are restricted to finite luminosity bins or flux-limited samples. Moreover, understanding any potential luminosity-dependent bias is crucial to probe the true distribution in the mass-luminosity plane.  \\end{itemize}  This paper is organized as follows. In \\S\\ref{sec:data} we describe the data; we present the traditional binned LF/BHMF in \\S\\ref{subsec:bin_BHMF} and describe the Bayesian approach in \\S\\ref{subsec:MCMC_BHMF}. We present our model results in \\S\\ref{sec:results}, discuss the results in \\S\\ref{sec:disc} and conclude in \\S\\ref{sec:con}. Throughout the paper we adopt a flat $\\Lambda$CDM cosmology with cosmological parameters $\\Omega_{\\Lambda}=0.7$, $\\Omega_0=0.3$, $h=0.7$, to match most of the recent quasar demographics work. Volume is in comoving units unless otherwise stated. We distinguish virial masses from true masses with a subscript $_{\\rm vir}$ or $_e$. Quasar luminosity is expressed in terms of the rest-frame 2500\\,\\AA\\ continuum luminosity ($L\\equiv \\lambda L_\\lambda$ or $l\\equiv \\log L$ for short), and we adopt a constant bolometric correction $C_{\\rm bol}=L_{\\rm bol}/L=5$. ",
        "conclusions": "\\label{sec:con}  This paper represents our first step towards a systematic investigation on the demographics of broad-line quasars in the mass-luminosity plane and its redshift evolution. We used a forward modeling Bayesian framework to model the joint distribution in the mass-luminosity plane. With simple model prescriptions for the underlying active BHMF and Eddington ratio distribution, we were able to fit the observed distribution above the sample flux limit and extrapolate below using constrained model parameters. We paid particular attention to the distinction between virial mass estimates and true masses, and corrected for the selection effects of flux limited samples. The main conclusions of the paper are the following:  \\begin{enumerate}  \\item[1.] Virial BH masses are not true masses (\\S\\ref{subsec:prelim}), and their explicit dependence on luminosity has implications for their conditional probability distributions: $p(m_e|m)\\neq p(m_e|m,l)\\neq p(m_e|l)$, and $\\sigma_{\\rm vir}>\\sigma_{lm}$. We found evidence that for \\MgII\\ and possibly \\CIV, there is a positive luminosity-dependent bias in virial masses, such that at fixed true mass, objects with luminosities above average have over-estimated virial masses. This is expected as there must be uncorrelated scatter between line width and luminosity at fixed true mass, due to the imperfectness of them being the surrogates for the virial velocity and the BLR radius in the virial mass technique. While this bias was noted earlier \\citep{Shen_etal_2008b,Shen_Kelly_2010}, this is for the first time that we quantified the level of this bias with the SDSS quasar sample. As a consequence, the dispersion in virial mass in narrow luminosity bins or flux-limited samples is generally smaller than the virial mass uncertainty \\citep{Shen_etal_2008b,Shen_Kelly_2010}, and the average virial BH masses in $z\\ga 0.7$ SDSS quasars are likely overestimated by a factor of a few.  \\item[2.] Our model LF is tightly constrained in the regime sampled by SDSS quasars, and makes reasonable predictions when extrapolated $\\sim 3$ magnitudes fainter, as compared with the LF measured from faint quasar surveys. Downsizing in LF evolution is clearly seen with our model LF. The break luminosity (as in a double power-law model) increases towards higher redshift. As a consequence, the LF slope for $M_i[z=2]<-27$ from a single power-law fit appears shallower at $z>3$ than at lower redshift.  \\item[3.] Except for a few \\texttt{zbins}, the model active BHMF usually cannot be well constrained to within a factor of a few. This is mainly caused by the additional free parameter on the luminosity-dependent bias, which leads to degeneracies with other parameters. The SDSS sample only probes the tip of the active SMBH population at high redshift, which makes it more difficult to constrain model parameters for the whole population. Nevertheless, the abundance of the most massive quasars is always overestimated using virial masses regardless of the luminosity-dependent bias. Within our model a turn-over at the low-mass end of the true BHMF is needed in order to fit the observed distribution in the mass-luminosity plane, but we caution that this result is based on our model assumptions and may suffer from the poor sampling of low-mass BHs in the SDSS sample. Further investigations are needed to test the robustness of this feature. This turn-over shifts to larger BH masses for the detected population in SDSS due to the flux limit. Although with large error bars, we found tentative evidence that downsizing also manifests itself in the active BHMF, and the number density of $M_{\\rm BH}\\sim 10^9\\,M_\\odot$ quasars peaks around $z\\sim 2$.  The total BH mass density in broad-line quasars is always insignificant compared with that in all SMBHs at any redshift (Fig.\\ \\ref{fig:soltan}).  \\item[4.] Within our model uncertainties, the Eddington ratio distribution at fixed true BH mass has a mean value that weakly depends on mass, and an average scatter of $\\sim 0.4$ dex. The sample-averaged Eddington ratio for all broad-line quasars ranges between $\\sim 0.01$ and $\\sim 0.3$, and increases with redshift (Fig.\\ \\ref{fig:eddington_dist}). The sample-averaged Eddington ratio for quasars above the SDSS flux limit is Malmquist-biased due to the scatter in Eddington ratios at fixed BH mass and the shape of the underlying active BHMF. Furthermore, using virial masses for the detected quasars tend to reduce the Eddington ratio again due to the luminosity-dependent bias (Fig.\\ \\ref{fig:eddington_dist}).  \\item[5.] Our model reproduces the observed distribution in the mass-luminosity plane (with virial masses). The flux limit, the scatter between true masses and virial masses, as well as the luminosity-dependent bias, change the distribution in the mass-luminosity plane significantly. Thus any features in the mass-luminosity plane based on virial masses must be interpreted with caution.   \\end{enumerate}  Our results highlight the need for understanding the systematics in the virial mass technique. All single-epoch virial mass estimators currently utilized in the literature are bootstrapped from the local RM AGN sample of only several dozen objects. These mass estimators have been extensively used for high redshift and high luminosity quasars, a regime that is poorly sampled by RM AGNs. We have shown that even without other systematic issues with the virial technique, the scatter and luminosity-dependent bias between virial masses and true masses already make it difficult to constrain the active BHMF accurately. Any additional systematic issues will simply make the situation worse.  Since BH mass is a crucial physical quantity and is related to many fundamental processes of SMBHs, refining the techniques that weigh the BH is of tremendous importance. As the most promising method to measure active BH masses, the RM technique (and its extensions of single-epoch virial estimators) should be tested and calibrated carefully. Progress is being made in reverberation mapping studies of broad-line AGNs and calibrations of virial mass estimators \\citep[e.g.,][]{Kaspi_etal_2007,Bentz_etal_2009b,Woo_etal_2010,Graham_etal_2011}, as well as in the statistical description of AGN variability as applied to reverberation mapping \\citep[e.g.,][]{Zu_etal_2011,Brewer_etal_2011,Pancoast_etal_2011}. However, to fully understand the BLR kinematics/geometry and to construct reliable mass estimators, a substantially larger RM sample is needed to sample an unbiased parameter space of broad-line objects, as well as a much better understanding of the systematics in the RM technique and single-epoch mass estimators."
    },
    "1107/1107.3849_arXiv.txt": {
        "abstract": "While the \\Swift~satellite is primarily designed to study gamma-ray bursts, its ultraviolet and optical imaging and spectroscopy capabilities are also being used for a variety of scientific programs. In this study, we use the UV/Optical Telescope (UVOT) instrument aboard \\Swift~to discover $0.5<z<2$ Lyman break galaxies (LBGs). UVOT has covered $\\sim266$ arcmin$^2$ at $>$60ks exposure time, achieving a limiting magnitude of $u<24.5$, in the Chandra Deep Field South (CDF-S). Applying UVOT near-ultraviolet color selection, we select 50 UV-dropouts from this UVOT CDF-S data. We match the selected sources with available multiwavelength data from GOODS-South, MUSYC, and COMBO-17 to characterize the spectral energy distributions for these galaxies and determine stellar masses, star formation rates (SFRs), and dust attenuations. We compare these properties for LBGs selected in this paper versus $z\\sim3$ LBGs and other CDF-S galaxies in the same redshift range ($0.5<z<2$), identified using photometric redshift techniques. The $z\\sim1$ LBGs have stellar masses of $\\langle \\rm Log~M_*/M_\\odot\\rangle = 9.4\\pm 0.6$, which is slightly lower than $z\\sim3$ LBGs ($\\langle \\rm Log~M_*/M_\\odot\\rangle=10.2\\pm 0.4$) and slightly higher compared to the $z\\sim1$ CDF-S galaxies ($\\langle \\rm Log~M_*/M_\\odot\\rangle = 8.7 \\pm 0.7$). Similarly, our sample of $z\\sim1$ LBGs has SFRs (derived using both ultraviolet and infrared data, where available) of $\\langle \\rm Log~SFR/(M_\\odot~yr^{-1})\\rangle=0.7 \\pm0.6 $, which is nearly an order of magnitude lower than $z\\sim 3$ LBGs ($\\langle \\rm Log~SFR/{\\rm M_\\odot~yr^{-1}}\\rangle=1.5 \\pm 0.4$), but slightly higher than the comparison $z\\sim1$ sample of CDF-S galaxies ($\\langle \\rm Log~SFR/{\\rm M_\\odot~yr^{-1}}\\rangle=0.2 \\pm 0.7$). We find that our $z\\sim 1$ UV-dropouts have $\\langle A_{\\rm FUV}\\rangle=2.0\\pm 1.0$, which is higher than $z\\sim3$ LBGs ($\\langle \\rm A_{\\rm FUV}\\rangle=1.0\\pm 0.5$), but is similar to the distribution of dust attenuations in the other CDF-S galaxies ($\\langle \\rm A_{\\rm FUV}\\rangle\\sim2.8\\pm 1.5$). Using the GOODS-South multiwavelength catalog of galaxies, we simulate a larger and fainter sample of LBGs to compare their properties with those of the UVOT-selected LBG sample. We conclude that UVOT can be useful for finding and studying the bright end of 0.5$<z<$2.0 LBGs. \\\\ ",
        "introduction": "\\label{sec:intro} Initial interest in UV-selected galaxies began as an attempt to find the most primeval galaxies \\citep[\\eg those that were theoretically predicted by][]{PP1}. As young systems, such galaxies are expected to have recent star formation, low metal enrichment and little dust. In addition to bright UV continua, young, star-forming galaxies are expected to have strong breaks at 912 \\AA, which occur as a result of the ground-state hydrogen ionization edge in the stellar absorption features of massive stars. The Lyman break technique exploits this feature in the rest-frame ultraviolet to isolate star-forming galaxies at great distances \\citep{Steidel92, Steidel93, Steidel95, Steidel00}, and the ``UV-dropout'' products of this technique are the so-called Lyman break galaxies (LBGs).  Despite having been selected as primordial systems, LBGs exhibit sufficient metal enrichment to exclude them from being the most primitive galaxies \\citetext{\\eg \\citealp{pettini02}\\ study the gravitationally lensed LBG, cB58, and find Type II supernovae residue, such as O, Mg, Si}. \\cite{Mori} have conducted high resolution hydrodynamic simulations that follow the chemical evolution of primordial galaxies, finding that LBGs resemble infant versions of elliptical and bulge systems in the local universe.  A decade of work has uncovered several significant results about LBGs at $z>2$ (see, e.g., \\citealt{giarev}\\ and references therein). LBGs form stars at intense rates, dominating the UV luminosity density at $z>2$. \\cite{Bouwens, Bouwens08, Bouwens09, Bouwens2010} find only a modest decrease in the UV luminosity density out to $z=6$, indicating that LBGs represent a major phase in the early stages of galaxy formation and evolution. UV-selected galaxies are valuable for mapping the evolution of the global star formation rate density \\citep{gia}, enriching the IGM \\citep{Adel03}, locating large-scale structure and quantifying galaxy environments \\citep{ouchi, adel05}.  While studying UV-selected galaxies at $z>1$ is valuable, measuring their physical properties is challenging since high redshift studies are biased toward observing more luminous galaxies, and faint features are difficult to detect at these great distances. At $z<1$, \\citet[further refined by \\citealt{Choopes}]{Heckman05}, employ far-UV (FUV) luminosity and surface brightness criteria to select LBG-analogs, thereby named Lyman break analogs (LBAs). These $z\\sim 0.2$ LBAs share several similar properties with LBGs: specific star formation rates (SFRs), metallicities, and attenuations \\citep{Heckman05, Choopes, me2}.  When artificially redshifted to $z\\sim3$, even their morphologies \\citep{rod, Overzier2010} and ionized gas kinematics \\citep{me-ifu, tsg} resemble those of actual LBGs. However, these galaxies are close enough to permit more detailed study of their physical properties.  Combined, these separate UV-selected samples (LBGs and LBAs) provide insight into the evolution of this important galaxy population. Furthermore, they can provide pertinent information about the cosmic star formation history. At $z>$1, UV-selected galaxies are plentiful. They contribute significantly to the total UV luminosity density at those redshifts \\citep{schimi04, Arnouts04}, but at $z<1$ these galaxies appear to be rare. Observations of galaxies at a multitude of wavelengths have shown that the star formation rate density (SFRD) of the Universe declines dramatically between $0 <z< 1$, peaking between $z\\sim1-3$ \\citep{Madau96, HB06}. The ``redshift desert'', named to signify the challenging nature of measuring redshifts for galaxies in the redshift regime $1<z<2$ \\citep{Renzini2009}, contains valuable information that connects the peak of star formation with its rapid decline. It is likely that in this redshift range, the Hubble sequence took shape \\citep{Papovich2005}.  A few studies have used UV and optical data to expand the LBG selection to intermediate redshifts: $1<z<2$, connecting the low redshift ($z<1$) LBA sample with the high redshift ($z>2$) LBG sample. \\cite{Ly2009} have targeted $1.5 <z<2.7$ LBGs, selected using the Galaxy Evolution Explorer \\citep[\\GALEX;][]{galex} near-UV (NUV) filter and deep Subaru optical filters; this work concludes that the peak of star formation occurred at $1.5<z<3$. \\cite{Burgarella2006} have selected LBGs using the \\GALEX~FUV$-$NUV color for galaxies in the COMBO-17 sample with known redshifts between $0.9 \\lesssim z \\lesssim 1.6$. This study compares infrared with ultraviolet observations and finds two populations: 24-micron detected and non-detected LBGs. The former can be classified as luminous infrared galaxies (LIRGs) with Log L$_{IR}>11$, exhibiting significant amounts of dust attenuation which is anti-correlated with the observed UV luminosity; the latter case appears to have little dust attenuation since stacking these LBGs constrains the infrared luminosity to be an order of magnitude less than the rest-frame NUV. Recently, two other studies have identified LBGs in this redshift range using Hubble Space Telescope (\\HST) WFC3 data \\citep{Hathi2010, Oesch2010} and have studied the evolution of luminosity function parameters with redshift. \\cite{Oesch2010} find that the faint-end slope of the luminosity function, $\\alpha$, appears to transition from a steep slope at $z>2$ to a flatter slope in the local universe.  Since the 1$<z<2$ regime is a pivotal period in galaxy evolution, we ask: How do the properties of galaxies at these redshifts compare to local star-forming galaxies or to LBGs at high redshifts? Deep observations of some fields, such as the Chandra Deep Field South (CDF-S), have detected numerous non UV-selected galaxies at these redshifts allowing us to investigate: How do UV-dropouts compare to other galaxies of the same redshifts?  In this paper we introduce another LBG sample that bridges the other aforementioned LBG samples across this relevant redshift range: 0.5 $\\lesssim z \\lesssim$ 2.  This sample of LBGs is selected using {\\it Swift} \\citep{Gehrels04} UV-Optical Telescope (UVOT) observations of the CDF-S. In this field, low redshift ($z \\leq$ 2) and high redshift ($z \\geq$ 3) LBGs have been identified \\citep{Burgarella2006, Hathi2010, Oesch2010, Vanzella2009, gia2004, Bouwens, Bouwens08, Bouwens09, Bouwens2010}, yet none using UVOT.  As a UV instrument, \\Swift~UVOT offers a unique method for selecting LBGs, complementary to using \\GALEX~and WFC3. \\GALEX, with the largest  field-of-view (FOV; 1$^\\circ$.2 diameter), covers the largest sky area of all these three UV instruments, yet is not as sensitive as WFC3 (\\ie \\GALEX~has imaged $\\sim 4700 \\deg^2$ down to 22.5 magnitudes as of December 2010\\footnote{Information available on \\GALEX~Legacy Survey website: {\\texttt http://galexgi.gsfc.nasa.gov/docs/galex/Documents/GALEX-Legacy-Survey.html}} and 80~deg$^{2}$ to 25.0 magnitudes in the Deep Imaging Survey). However, UVOT offers two advantages over \\GALEX: multiple NUV filters (\\GALEX~only has one wide NUV filter), which allows the selection of UV-dropout candidates based on UVOT data alone, and better spatial resolution (FWHM$\\approx$2\\arcsec.7, compared to FWHM$\\approx$5\\arcsec~for \\GALEX). WFC3 has excellent spatial resolution (0\\arcsec.2) and sensitivity ($\\geq$500 times that of UVOT; see Figure \\ref{fig:filters}); it also has multiple UV filters for UV-dropout selection. However, the UVOT FOV (17\\arcmin$\\times$17\\arcmin) is significantly larger than WFC3's FOV (2.7\\arcmin$\\times$2.7\\arcmin) and \\Swift~observations have covered a large area of the sky (\\eg there are 40 GRB fields with $\\approx$200 ks of exposure time in the UV filters). We can use UVOT observations to select these rare objects to study the bright end of the LBG population. Here we discuss the utility of using \\Swift~UVOT to select LBGs from the ``redshift desert''.  The CDF-S region of the sky has been extensively and deeply covered by multiwavelength observations. We use available datasets in several ways: to constrain various properties of the UVOT-selected LBGs; to provide a large comparison sample of other photometrically determined $0.5<z<2.0$ galaxies; and to draw from this large parent population a sample of simulated UV-dropout galaxies to compare their properties with the observed UVOT-selected LBGs. In Section \\ref{sec:data}, we present the sample selection and data analysis. We discuss the LBG candidates in Section \\ref{sec:results} and present our results for the $0.5<z<2.0$ LBG sample and compare these LBGs with the other $0.5<z<2.0$ galaxies (including a simulated LBG sample, defined in Section \\ref{sec:data}) and other UV-selected populations (LBAs and $z\\sim$3 LBGs). In Section \\ref{sec:end} we summarize our main results and discuss the merit of studying intermediate redshift ($0.5<z<2.0$) LBGs with \\Swift. Throughout our analysis, we use the following cosmology: H$_0$=70 km s$^{-1}$ Mpc $^{-1}$, $\\Omega_m$=0.30, and $\\Omega_\\Lambda$= 0.70, and assume a \\cite{Kroupa} initial mass function (IMF). ",
        "conclusions": "\\label{sec:end} In this paper, we used \\Swift~UVOT to identify and select LBG candidates. Based on redshifts determined from other multiwavelength catalogs (\\eg MUSYC, COMBO-17 and ACS-GOODs), we determined that 50 candidates were LBGs. We use these multiwavelength catalogs to fit SEDs to determine stellar masses, k-corrected absolute magnitudes and star formation rates. Using available MIPS data for 40 of the galaxies, we also determined UV$+$IR SFRs and FUV dust attenuations, A$_{\\rm FUV}$. GEMS and ACS-GOODS images offer high resolution optical morphologies of these galaxies. From our study we have determined that \\Swift~UVOT can select 0.5$<z<$ 2.0 LBGs within deep observed fields, but redshift confirmation or further photometric data is required as we found contamination (in $\\sim10-15$\\% of the candidates) from low redshift interlopers.  We found that the UVOT-selected LBGs have similar morphologies to $z=0.2$ LBAs from \\cite{rod}, $z=1.5$ emission line selected star-forming galaxies, and $z=4$ LBGs from \\cite{Lotz06} (Figure \\ref{fig:morph}). These LBGs also have similar values for SFR, stellar mass and FUV dust attenuation, but span larger ranges, compared with $z\\sim3$ LBGs and $z\\sim0.2$ LBAs. However, compared to the $z\\sim1$ ACS sample, these LBGs are bluer and brighter and have slightly higher stellar masses, and marginally higher total SFRs.  We used best-fit SEDs to simulate UVOT photometry for the $z\\sim1$ ACS sample and selected a sample of simulated LBGs based on the same criteria applied on the observed UVOT-selected LBGS. Despite including fainter [$\\langle {\\rm z_{850} (simulated)} \\rangle \\approx \\langle {\\rm z_{850} (observed)} \\rangle +2.5$ mag] galaxies, the simulated sample does exhibit similar properties as the observed UVOT-selected LBG sample: slightly more massive, more UV-luminous, and slightly higher SFRs compared to the $z\\sim1$ ACS sample (Figures \\ref{fig:dist} and \\ref{fig:cmd}). The LBGs (observed and simulated) have similar dust attenuation compared to the $z\\sim1$ ACS comparison sample and the amount of FUV attenuation is not as low as in the $z\\sim3$ LBGs or $z\\sim0.2$ LBAs. We find that  $\\sim32\\%$ of the IR-detected LBGs are LIRGs. This fraction is not as high as what \\cite{Burgarella2006} found for their $z\\sim1$ \\GALEX-selected LBGs, but it is a significant fraction of the LBGs. Red (U$-$I$>2$) galaxies were found in both the simulated LBG sample ($\\sim 3$\\%) and the observed LBG sample (2 galaxies, or $\\sim4$\\%).  This research investigates the viability of using \\Swift~UVOT to select intermediate redshift ($z=0.5 - 2$) LBGs. Until recently, there was little opportunity to study LBGs in this redshift desert. However, with the WFC3 upgrades to \\HST~the new ultraviolet filters have been used to efficiently select hundreds of LBGs in this regime \\citep{Hathi2010, Oesch2010}. While WFC3 has higher sensitivity and can select fainter LBGs more efficiently, its FOV is small; GALEX data has larger sky coverage, yet it suffers source confusion from low spatial resolution and does not have multiple NUV filters. We have shown that UVOT can be valuable for studying the bright end of the LBG sample. Having compared with a simulated LBG sample, we learn that the fainter, simulated LBGs appear to resemble the observed, brighter LBGs in important properties (\\eg stellar mass, SFR, dust attenuation). According to statistical studies of the contribution of UV-luminous galaxies to the SFR density \\citep{schimi04}, the number of LBGs decreases rapidly in the low redshift Universe ($z<2$). Since UVOT has covered a larger fraction of the sky than WFC3, UVOT data has the potential to recover a larger (albeit, shallow) sample of LBGs in this redshift range. There are several other deep UVOT fields (40 GRB fields with UV exposure times of $\\geq$ 200 ks) for which a similar analysis can be done to further study LBGs between $0.5<z<2$.  As future missions, such as James Webb Space Telescope (JWST), push to discover more distant (rest-frame UV-selected) galaxies, it becomes more important to understand the selection of these galaxies in the relatively nearby Universe, as well as how their properties compare with the higher redshift ($z>2$) population."
    },
    "1107/1107.3552_arXiv.txt": {
        "abstract": "We present the discovery and analysis of two ultra-luminous supernovae (SNe) at z $\\approx0.9$ with the Pan-STARRS1 Medium-Deep Survey. These SNe, PS1-10ky and PS1-10awh, are amongst the most luminous SNe ever discovered, comparable to the unusual transient SCP 06F6. Like SCP 06F6, they show characteristic high luminosities ($M_{\\rm bol} \\approx -22.5$ mag), blue spectra with a few broad absorption lines, and no evidence for H or He. We have constructed a full multi-color light curve sensitive to the peak of the spectral energy distribution in the rest-frame ultraviolet, and we have obtained time-series spectroscopy for these SNe. Given the similarities between the SNe, we combine their light curves to estimate a total radiated energy over the course of explosion of $(0.9-1.4) \\times 10^{51}$ erg. We find expansion velocities of $12,000-18,000$ km s$^{-1}$ with no evidence for deceleration measured $\\sim$3 rest-frame weeks either side of light-curve peak, consistent with the expansion of an optically-thick massive shell of material. We show that radioactive decay is not sufficient to power PS1-10ky, and discuss two plausible origins for these events: the initial spin-down of a newborn magnetar in a core-collapse SN, or SN shock breakout from the dense circumstellar wind surrounding a Wolf-Rayet star. ",
        "introduction": "The observational and physical parameter space occupied by supernovae (SNe) has expanded dramatically because of the recent discovery of several ultra-luminous SNe. These SNe are significantly more luminous than Type Ia explosions, displaying absolute bolometric magnitudes at light curve maximum of $< -21$ mag and total radiated energies on the order of $10^{51}$ erg. Thus far, ultra-luminous SNe have displayed impressive diversity, ranging from the Type Ic SN 2007bi, proposed to be a pair-instability explosion \\citep{GalYam_etal09, Young_etal10}, to the Type IIn SN 2006gy, with strong signs of circumstellar interaction \\citep{Smith_etal07, Smith_McCray07, Ofek_etal07}. The mysterious transient SCP 06F6 discovered by \\citet{Barbary_etal09} showed a largely featureless spectrum with a few broad absorption lines, and proved so perplexing that no redshift could be identified. Because it was uncertain if the source was Galactic or at a cosmological distance, many ideas were presented to explain it ranging from an outburst on a white dwarf to a broad-lined QSO \\citep{Barbary_etal09, Gansicke_etal09, Chatzopoulos_etal09, Soker_etal10}.  Since SCP 06F6 was discovered by \\citet{Barbary_etal09}, similar objects have been found with significant frequency in wide-field transient surveys. \\citet{Quimby_etal11} collected five additional SCP 06F6-like sources discovered by the Texas Supernova Survey and the Palomar Transient Factory (PTF), including the earlier-identified SN 2005ap \\citep{Quimby_etal07}. They identifed narrow \\ion{Mg}{2} $\\lambda\\lambda$2796,2803 absorption lines in their spectra, presumably associated with the interstellar media of the host galaxies. From this absorption, they derived redshifts ranging from $0.25 \\lesssim \\textrm{z} \\lesssim 0.5$ for their sources, and a redshift of z $= 1.19$ for SCP 06F6 itself. These cosmological distances imply that the SCP 06F6-like sources are some of the most luminous SNe known, with peak absolute magnitudes $M_{U} \\lesssim -22$ mag and total radiated energies $\\gtrsim 10^{51}$ erg.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{f1} \\caption{Cut-outs of PS1 \\ips-band images showing the region around PS1-10awh (top) and PS1-10ky (bottom). The left column shows stacked images using data before explosion; the right column shows images from single nights around maximum light.} \\label{comparison} \\end{figure}  With their redshifts known, a common set of observational properties began to emerge for the SCP 06F6-like objects. In addition to their very high peak luminosities, these sources all show distinctive symmetric light curves with rise and decline times on the order of 30 days in the rest frame. They also all show blue spectra with only a handful of features; \\citet{Quimby_etal11} identify the broad absorption lines with light metals (C, O, Si, and Mg). There is no clear evidence for H or He in the spectra of an SCP 06F6-like transient, although \\citet{Quimby_etal07} note a broad feature in SN 2005ap that may be associated with H$\\alpha$. There also is no detection of strong \\ion{Fe}{2} or \\ion{Fe}{3} lines in the early or peak spectra of these transients.  Subsequently, the ultra-luminous SN 2010gx was discovered by both the Catalina Real-time Transient Survey \\citep{Mahabal_etal10, Mahabal_Drake10} and PTF \\citep{Quimby_etal10}, and was determined to be at z = 0.23. \\citet{Pastorello_etal10} carried out a detailed study of this source, which shows SCP 06F6-like spectral features and high luminosity at early times. They find that, a few weeks after peak light, this source begins to show iron and other features indicative of Type Ic SNe. They assert that SCP 06F6-like objects are likely to be associated with SNe Ibc, consistent with the lack of H in their spectra.  Two potential explanations for these SCP 06F6-like sources have surfaced: the interaction of the SN shock with a dense circumstellar shell of H-poor material \\citep[e.g.,][]{Chevalier_Irwin11} or the spin-down of a new-born magnetar embedded in the SN ejecta \\citep{Kasen_Bildsten10, Woosley10}. Both of these models have the potential to explain the high luminosities and association with SNe Ibc.  Here, we present two new ultra-luminous SNe at z $\\approx$ 0.9 discovered with Panoramic Survey Telescope \\& Rapid Response System 1 (Pan-STARRS1, abbreviated as PS1 here). By combining data for these two sources, we obtain the first full composite multi-color light curve for SCP 06F6-like objects and place constraints on the origins of these sources. Our photometry in multiple filters allows us to measure the color evolution of these sources, while their high redshifts aid in measuring bolometric luminosities, as these sources' spectral energy distributions (SEDs) peak in the rest-frame ultraviolet. Our spectroscopic observations sample the entire light curve, enabling us to measure the evolution of photospheric velocities. These measurements place important constraints on any model that attempts to explain these sources' high luminosities.  In \\S \\ref{obs}, we describe our data, obtained using PS1, MMT, Gemini, \\emph{GALEX}, and the EVLA. In \\S \\ref{host} we constrain the properties of the host galaxies. In \\S \\ref{temp}, \\ref{lum}, and \\ref{vel}, we measure the evolution of their photospheric temperatures, bolometric luminosities, and photospheric velocities, combining our new data with published data of SCP 06F6 and SN 2010gx. In \\S \\ref{model}, we consider three possible physical scenarios that might explain the characteristics of these sources: radioactive decay, magnetar spin-down, and circumstellar interaction. Finally in \\S \\ref{conclude}, we summarize our results. ",
        "conclusions": "\\label{conclude}  Using multi-color photometry and multi-epoch spectroscopy, we find that SCP 06F6-like objects radiate $\\sim10^{51}$ erg in just a few months, making them amongst the most energetic SNe known. Their peak bolometric magnitudes are typically $\\lesssim -$22.5 mag, as compared with $-$19.5 mag for SNe Type Ia \\citep{Contardo_etal00}, $-$21.4 mag for the proposed pair-instability SN 2007bi \\citep{GalYam_etal09, Young_etal10}, and $-$21.8 mag for the ultra-luminous Type IIn SN 2006gy \\citep{Smith_etal07, Ofek_etal07}. Our estimates are consistent with the lower-redshift sources presented in \\citet{Quimby_etal11}.  There is no evidence for decreasing radial velocity during the SCP 06F6-like phase in the objects studied here, consistent with the expansion of an optically-thick shell predicted by both the magnetar and circumstellar interaction scenarios. However, the SCP 06F6-like SN 2005ap provides a counter-example to this point, displaying a clear decrease in velocity \\citep{Quimby_etal07}. A larger sample of objects with time-series spectroscopy is required to test the prevalence of decelerating photospheric velocities in SCP 06F6-like objects. One source to date, SN 2010gx \\citep{Pastorello_etal10} has shown evidence for dissipation of the optically-thick shell and a decrease in optical depth, revealing a normal SN Type Ic in its interior.  Here, we show that radioactive decay can not explain SCP 06F6-like objects, which rules out pair-instability models \\citep[e.g.,][]{Barkat_etal67} for this particular class of ultra-luminous SNe, unlike in the case of the unusually luminous SN 2007bi \\citep{GalYam_etal09}. We investigate two physical scenarios that both require an energetic optically-thick shell: the spin-down of a new-born magnetar or shock breakout from a dense circumstellar medium. Both scenarios can fit our data with plausible parameters for a SN Type Ic, although both require rather extreme additional conditions. If magnetars are responsible for SCP 06F6-like objects, they need to be spinning near breakup ($1-2$ ms period). Meanwhile, if shock breakout causes these events, the progenitor must have undergone an LBV-like outburst, expelling several solar masses of H-poor material in just a few years before stellar death; however, no such violent eruption has ever been observed from the Wolf-Rayet stars that are likely the progenitors of these SNe (although dense He-rich CSM has been observed to be present around some SNe Type Ib; \\citealt{Pastorello_etal07, Foley_etal07}). One of the most promising strategies for distinguishing between these scenarios in the future is early-time measurement of the photospheric radius, because at the time of light curve rise in the CSM interaction model, the SN will expand from a relatively much larger radius as compared with the magnetar spin-down model. In addition, late-time observations have the potential to distinguish between the two models. For example, the later observations of SN 2010gx may conflict with expectations from the magnetar model, as this scenario predicts that the SN ejecta themselves are being swept into a thin shell by the magnetar wind; however, relatively normal SN Type Ic ejecta are observed in this source at late times.  Any model for SCP 06F6-like sources needs to explain why these ultra-luminous events are only associated with a small fraction of SN explosions; the extreme conditions implied by both models discussed here might justify their rarity, but it remains to be determined if we can expect such fast-spinning magnetars or dense H-poor circumstellar media in sufficient numbers to explain the rates of SCP 06F6-like sources. The lack of detection of host galaxies for PS1-10ky and PS1-10awh (down to $0.02\\ L^{\\star}$) implies that they are in dwarf galaxies with low metallicities and high specific SFRs. This is a common feature for these type of transients \\citep{Quimby_etal11, Pastorello_etal10, Neill_etal11}. As none have been found in an $L^{\\star}$-type galaxy, this may point to metallicity playing a key role in the progenitor channel---or it may simply be an artifact of the bulk of cosmic star formation occuring in systems with high specific SFR. A larger sample of SCP 06F6-like sources is required to distinguish between these scenarios.  Moreover, we note that these sources display some degree of homogeneity as a class, showing peak luminosities, expansion velocities, and temperature evolution which agree to better than a factor of two across our (albeit small) sample. Given the widespread cut-off in luminosity blueward of 2500 \\AA, we believe SCP 06F6-like sources will be of limited utility for ultraviolet high-redshift absorption-line studies (e.g., Ly $\\alpha$), in contrast with the claims of \\citet{Quimby_etal11}. As wide-field surveys like PS1 and PTF continue to operate, many more SCP 06F6-like sources will be discovered; PS1 is currently discovering $\\sim$1 such source each month at z $\\gtrsim$ 0.5. We will be able to test their diversity as a class and measure their rates. Increased sample sizes at a range of redshifts will better constrain the physical properties of these highly energetic explosions, their host stellar populations, and discern between the magnetar spin-down and shock breakout scenarios."
    },
    "1107/1107.4833_arXiv.txt": {
        "abstract": "We argue that the bright flare of the binary pulsar \\object{PSR~B1259$-$63/LS2883} detected by the {\\it Fermi} Large Area Telescope (LAT), is due to the inverse Compton (IC) scattering of the unshocked electron-positron pulsar wind with a Lorentz factor $\\Gamma_0 \\approx 10^4$.  The combination of two effects both linked to the circumstellar disk (CD), is a key element in the proposed model. The first effect is related to the impact of the surrounding medium on the termination of the pulsar wind.  Inside the disk, the ``early'' termination of the wind results in suppression of its gamma-ray luminosity.  When the pulsar escapes the disk, the conditions for termination of the wind undergo significant changes. This would lead to a dramatic increase of the pulsar wind zone, and thus to the proportional increase of the gamma-ray flux. On the other hand, if the parts of the CD disturbed by the pulsar can supply infrared photons of density high enough for efficient Comptonization of the wind, almost the entire kinetic energy of the pulsar wind would be converted to radiation, thus the gamma-ray luminosity of the wind could approach to the level of the pulsar's spin-down luminosity as reported by the {\\it Fermi} collaboration. ",
        "introduction": "\\label{sec:intro}  The pulsar magnetospheres are effective gamma-ray emitters \\citep{fermi_pulsars}. Pulsars can produce potentially detectable gamma-ray emission also due to the bulk Comptonization of their cold ultrarelativistic winds. While in the case of isolated pulsars this radiation component is generally weak, unless the wind is accelerated (relatively) close to the pulsar \\citep{bogovalov00,aharonian12}, in binary systems the radiation is significantly enhanced due to the presence of dense target photon field supplied by the optical companion \\citep{ball00,ball01,khangulyan07,khangulyan11}.  Gamma-rays are produced also after termination of the wind due to the IC scattering of relativistic electrons accelerated by the termination shock.  In isolated pulsars the termination of the wind leads to the formation of the so-called Pulsar Wind Nebulae with a distinct extended synchrotron and IC emission which can be easily separated from the radiation of the unshocked wind. In binary pulsars, the separation of the radiation components produced before and after the termination of the pulsar wind is a more difficult task; it requires careful analysis of the spectral and temporal features of two radiation components.    The most promising object for exploration of processes of formation, acceleration and termination of pulsar winds in binary systems is \\object{PSR~B1259$-$63/LS2883}.  It contains a 47.7~ms pulsar orbiting a luminous star in a very eccentric orbit  \\citep[eccentricity $e=0.87$, period $P_{{\\rm orb}}= 1237$~d, semi-major axis $a_{\\rm   2}=7.2\\rm \\, AU$, see][and references therein]{johnston92,negueruela10}.  The variable TeV gamma-ray emission detected from this object by the {\\rm H.E.S.S.}  collaboration \\citep{aharonian05,aharonian09}  was in general agreement  with the  predictions  \\citep{kirk99}, but the TeV gamma-ray lightcurve appeared to be  different from expectations. Several possible scenarios suggested for interpretation of the reported TeV  lightcurve \\citep{khangulyan07}, predict essentially different gamma-ray fluxes at  MeV/GeV energies, therefore the measurements in this energy band should significantly contribute to the understanding of the hydrodynamics and the particle acceleration processes in binary pulsars.  The recent observations of \\object{PSR~B1259$-$63/LS2883} with {\\it   Fermi} LAT revealed a complex behavior of the system in GeV gamma-rays \\citep{abdo11,tam11}. Close to the periastron passage, the source has been detected with a modest  energy flux of $6\\times10^{-11}\\rm erg\\,cm^{-2}\\,s^{-1}$. However, $+30$ days after periastron passage (with largest recorded flux level at $+35$ days),  a spectacular flare has been recorded.  It lasted approximately two weeks with an enhanced flux above 100 MeV at the level of $3\\times10^{-10}\\rm erg\\,cm^{-2}\\,s^{-1}$ \\citep{abdo11,tam11}.  The flare is  characterized by a sharp increase and a rather smooth decay over approximately $2$ weeks. Importantly, the strong increase of gamma ray flux was not accompanied by a noticeable  change of the X-ray flux \\citep{abdo11}. While the  weak GeV flux reported during the periastron passage can be explained by IC scattering of electrons accelerated at the termination shock  and/or  by the bulk Comptonization of the unshocked electron-positron wind \\citep{khangulyan11}, the post-periastron flare was a real surprise.  In general,  this flare  represents a unique case in astrophysics when the available energy of an object is fully  converted to nonthermal high energy radiation. On the other hand, given the current belief  that the rotational energy of the pulsar is converted to kinetic energy of cold ultrarelativistic electron-positron wind, the binary pulsars can in principle operate as perfect `100 \\%' efficient gamma-ray emitters if the density of the surrounding target photon field would be high enough  for realization of Comptonization of the pulsar wind in the saturation regime. Yet, any explanation of this flare should address several important issues, namely (i) the extraordinary high luminosity of gamma radiation at the level close to the pulsar's spin-down luminosity (SDL), $\\dot E=8 \\times 10^{35} \\ \\rm erg/s$; (ii)  lack of  enhanced non-thermal X-ray flux;  (iii) orbital phase of the flare. ",
        "conclusions": "In this letter we propose a model for the recent {\\it Fermi} observations of gamma-rays from \\object{PSR~B1259$-$63/LS~2883}. We assume that gamma-rays are produced at interaction of the cold pulsar wind with external photon fields. While the Comptonization of the wind by radiation of the companion star can explain the modest  gamma-ray flux detected during the periastron passage, the order of magnitude higher flux of the post-periastron flare requires an additional seed photon component.  We propose that the radiation of the heated CD can play this role.  The Comptonization of the pulsar wind by the CD emission proceeds in a quite specific way.  Inside the disk it is  suppressed  because of the high ram pressure in the disk. The density of the stellar optical photons generally dominates over the density of IR photons. But in the immediate vicinity of the disk, where the IR density can be  still very high, the optical depth for Compton scattering might  be significantly larger than 1.  This would lead to a very effective, close to 100\\% efficiency of transformation of the rotational energy of the pulsar to high energy gamma rays via Comptonization of the cold ultrarelativistic wind. In the suggested scenario, the departure of the pulsar wind from the disk will lead to a gradual decrease of the photon density from the disk and correspondingly to the reduction of the IC gamma-ray flux on time-scales of days.   During the exit of the pulsar from the disk at the epoch  before the periastron passage, the termination shock is expected to expand towards the direction opposite to the observer, thus we should not expect strong gamma-ray signal. This agrees with the {\\it Fermi LAT} observations. The main issue for realization of the proposed scenario is whether the pulsar-disk interaction can provide the required photon field.  This is a key question which should be answered by detailed studies of interaction of the pulsar wind with the CD. This study should also address another important issue, which is related to the 15 days delay between the appearance of the pulsed radio emission and the epoch of the GeV flare.  If the proposed interpretation is correct,  this would be the second case of direct measurement of  a pulsar wind's  Lorentz factor. Recently,  the Lorentz factor of order of $10^6$ has been  extracted from the observation of very high energy pulsed gamma-ray emission from the Crab Pulsar \\citep{aharonian12}. Interestingly,  the two orders of magnitude smaller Lorentz factor of the pulsar wind  in \\object{PSR~B1259$-$63/LS~2883} compared to the Crab pulsar, is quite close  the ratio of spin-down luminosities of two pulsars. If this is not just a coincidence, the linear dependence $\\Gamma_0 \\propto L_{\\rm SD}$  would imply a universal production rate of electron-positron pairs of order of $10^{38} \\ \\rm e^\\pm /sec$,   independent of the pulsar spin-down luminosity.   Finally, we should note that in principle the gamma-ray flare can be explained also by the post-shock flow. Since the flux detected by {\\it Fermi} LAT was close to pulsar's SDL, this interpretation requires significant Doppler boosting of radiation. Relativistic flows can be indeed formed at interactions of the pulsar and SWs \\citep{bogovalov08}, and therefore the broad-band radiation can be affected by the Doppler boosting \\citep{khangulyan08,dubus10}.  The Doppler boosting should amplify also the X- and TeV gamma ray fluxes.  However, the lack of any noticeable activity during the flare at other wavelengths \\citep{abdo11} makes this interpretation less likely.  Also, one should note that the flux level of the Doppler boosted radiation can easily exceed the limit set by the pulsar's SDL, but flares with apparent super SDL so far have not been detected.  On the other hand, the detection of such events in future observations would rule out the origin of gamma-radiation related to the unshocked pulsar wind."
    },
    "1107/1107.5984.txt": {
        "abstract": "Gravitational waves coming from Super Massive Black Hole Binaries (SMBHBs) are targeted by both Pulsar Timing Array (PTA) and Space Laser Interferometry (SLI). The possibility of a { single} SMBHB being tracked first by PTA, through inspiral, and later by SLI, up to { merger}  and { ring down}, has been previously suggested. Although the bounding parameters are drawn by the current PTA or the upcoming Square Kilometer Array (SKA), and by the New Gravitational Observatory (NGO), derived from the Laser Interferometer Space Antenna (LISA), { this paper also addresses} sequential detection beyond specific project constraints. We consider PTA-SKA, { which is} sensitive from $10^{-9}$ to $p\\cdot10^{-7}$ Hz ($p=4,~~8$), and SLI, { which operates} from $s\\cdot10^{-5}$ up to $1$ Hz ($s = 1,~~3 $). A SMBHB in the range $2\\cdot 10^8 - 2 \\cdot 10^9~M_\\odot$ (the masses are normalised to a $(1+z)$ factor, { the red shift lying between  $z = 0.2$ and $z=1.5$}) moves from the PTA-SKA to the SLI { band over a period} ranging from two months to fifty years. By combining three Super Massive Black Hole (SMBH)-host relations with three accretion prescriptions, nine astrophysical scenarios are formed. They are then related to three levels of pulsar timing residuals ($50$, $5$, $1$ ns), generating twenty-seven cases. For residuals of $1$ ns, sequential detection probability will never be better than $4.7 \\cdot 10^{-4}$ y$^{-2}$ or $3.3 \\cdot 10^{-6}$ y$^{-2}$(per year to { merger}  and per year of survey), according to { the} best and worst astrophysical scenarios, respectively; { put differently this means} one sequential detection every $46$ or $550$ years for an equivalent { maximum} time to { merger} and duration of the survey. The chances of sequential detection are further reduced by increasing values of the $s$ parameter (they vanish for $s = 10$) and of the SLI noise, and by decreasing values of the remnant spin. The spread in the predictions diminishes when { timing precision is improved or the SLI low frequency cut-off is lowered}.  So while transit times and the SLI Signal to Noise Ratio (SNR) may be adequate, the likelihood of sequential detection is { severely hampered} by the current estimates on the number - just an handful - of { individual inspirals observable} by PTA-SKA, and to a lesser extent by the { wide gap} between the pulsar timing and space interferometry bands, and by the severe requirements on pulsar timing residuals. Optimisation of future operational scenarios for SKA and SLI is briefly dealt with, since a detection of even a single event would be of paramount importance for the understanding of SMBHBs and of the astrophysical processes connected to their formation and evolution. ",
        "introduction": "At different paces, Earth-based and space gravitational wave detectors and observatories are already - or will { soon} be - operating. They { aim} to cover different parts of the gravitational spectrum, targeting an extremely large variety of astrophysical sources. Nonetheless, large gaps will show between the sensitive frequency bands. One such gap { is that} between the Pulsar Timing Array (PTA) - in the future the Square Kilometer Array (SKA) - and the Space Laser Interferometry (SLI) bands. SLI may materialise as the New Gravitational Observatory (NGO), { derived from} the Laser Interferometer Space Antenna (LISA). Though their bands lie apart, PTA-SKA and SLI may monitor the same type of source, namely a Super Massive Black Hole Binary (SMBHB) at different evolutionary stages. In this paper, we explore under what circumstances { a single} SMBHB may be viewed first by PTA-SKA and later by SLI.  { The motivation for the study lies} in the opportunities offered by the analysis of the same sources with different instruments ({ radio astronomy} and interferometry), and at different relativistic regimes (inspiral, coalescence, { merger}  and ring-down). If pulsar timing { were to} provide mass and spin parameters, the latter could be counter-checked by laser interferometry. Likewise, a binary system could be examined, {\\it vis--\\`a--vis} the presence of matter, gas and other bodies, at different phases of { its} evolution.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\subsection{{ Observation} by pulsar timing}  In the frequency band between $10^{-9}$ Hz and some fraction of $10^{-6}$ Hz, PTA offers the unique chance to observe gravitational radiation. Beyond the { observation} of a stochastic background, the { challenge of} { observing a single} SMBHB by PTA has recently received growing attention. Simulations { concur} { in predicting } a scenario wherein some SMBHBs stick out of the stochastic { signal} floor of unresolved SMBHBs. { The observation of an individual SMBHB} provides opportunities for { new measurements} in general relativity { and new perspectives for the scientific community}.  { The state} of the art on timing residual precision\\footnote{The timing residuals are computed from the phase difference between the observed time of arrival (ToA) and the predicted ToA, based on the current model parameters.} from current  PTAs\\footnote{PPTA {\\scriptsize http://www.atnf.csiro.au/research/pulsar/ppta}, {\\scriptsize EPTA http://www.epta.eu.org/}, NANOGrav{\\scriptsize http://nanograv.org/}.} { lies} in the $100-50$ ns domain{ , while} improvement by SKA\\footnote{SKA {\\scriptsize http://www.skatelescope.org/}.} down to $10-1$ ns is expected (Liu et al. 2011).  Investigations\\footnote{In the following cited studies, two simplifying { hypotheses} have been adopted: i) binaries are on circular orbits; ii) the mergers are gravitational wave driven.} on individual sources aim to recover { physical parameters such as the spins and the masses} of, and the distance to, an SMBHB (Sesana \\& Vecchio 2010) using the gravitational { wave front} curvature (Deng \\& Finn 2011), or the Pulsar term\\footnote{The 'Earth term' is the gravitational wave strain at the Earth at the time when the pulse is received. The 'Pulsar term' is the strain at the pulsar at the time when the pulse is emitted. An SMBHB produces two quasi-monochromatic components in PTA residuals, and likely of different frequencies as the SMBHB evolves.} (Corbin \\& Cornish 2010; Lee et al. 2011; Mingarelli et al. 2012). From these studies it emerges that the distance to a pulsar is important for individual { SMBHBs}, { unlike} statistical background observation (Jenet et al. 2004; Sesana et al. 2009; Finn \\& Lommen 2010). { Finn and Lommen (2010) and Pitkin (2012) analysed} bursts coming from different sources including individual { SMBHBs}.  Burt et al. (2011) suggest { refraining from searching for} new pulsars, unless close to an { existing cluster} of good pulsars; instead, they recommend { the allocation} of more { observation} time to already low-noise pulsars.  Turning to specific { observation} targets, searches have so far produced negative results: Jenet et al. (2004) { found no evidence of the emission} of gravitational waves by a { supposed} SMBHB in 3C 66B; { neither did} Lommen \\& Backer (2001), who were seeking evidence for an SMBHB in Sgr A*.  Finally, Yardley et al. (2010) describe the { observations} used to produce the sensitivity curves for the Parkes { { radio telescope}}, and propose a method for detecting significant sinusoids in { PTA}.  \\subsection{Detection by space laser interferometry}  Gravitational wave detection in space was first proposed by means of a small sized interferometer on-board a single satellite (Grassi-Strini et al. 1979), before shaping into a triangular satellite configuration (Bertotti 1984; Faller et al. 1985,1989). { The} LISA Pathfinder (Antonucci et al. 2012) is deemed an important step { towards technological} maturity.  { NGO\\footnote{\\scriptsize http://sci.esa.int/ngo} (ESA 2011), which { is derived from the previous LISA proposal}\\footnote{\\scriptsize http://sci.esa.int/lisa}, is a space project} designed to measure gravitational radiation over a broad band at low frequencies, where the Universe is richly populated by strong sources of gravitational waves, including SMBHBs. NGO plans to trace the formation, growth and merger history of SMBHs during different epochs, % (growth of quasi-stellar objects and widespread star formation, or when the universe was less than $1$ Gyr old) measuring spin and masses, with an unprecedented precision, often where the Universe is blind with our current electromagnetic techniques (ESA 2011). %, often where the Universe is blind with our current electromagnetic techniques (ESA 2011). In fundamental physics, different tests on general relativity, including the no-hair theorem and the dynamics in strong-field, and on alternative theories, will be feasible with { SLI}. NGO { (Amaro--Seoane et al. 2012a,b)} implies a shift to higher frequencies of the sensitivity band, as compared to LISA.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\subsection{Sequential detection}  Pitkin et al. (2008) { first proposed} sequential detection, but { it appeared} necessary to improve and update   their initial work for several reasons. { We begin by ascertaining} that the total SMBHB mass is generally larger than the fifty million solar masses considered by Pitkin et al. (2008).  New { analyses} are carried out herein. First, we examine various astrophysical scenarios { combining} SMBH-host relations and accretion processes, { and} dry and wet mergers. Second, we span a large range of residuals ($50$--$1$ ns). Third, we estimate the number of events and the probability of sequential detection, building our investigation upon the recent { statistical} findings of individual detection by PTA-SKA, { which were} not available at the time of the work { by} Pitkin et al. (2008). Fourth, we present the Signal to Noise Ratio (SNR) of SLI for the sources concerned.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\subsection{Structure of the paper}  Section 2 is devoted to the computation of transit times (how long it takes for a binary to switch { bands}) from the PTA-SKA to the SLI band for a range of SMBHB masses. In the same section, the ringing frequencies are computed for { the} (stationary and rotating) { SMBHB remnants}.  Section 3 is the core of the { paper, in which we} study the impact of astrophysical, observational and detection constraints separately, one step after the other, steadily building up our analysis. Details of the models are described by Sesana et al. (2009), to which { the reader is referred for further information}. Super Massive Black Hole (SMBH)-host relation models { are then combined} with accretion prescription models, and different values of timing residuals { are considered}. We determine first the { maximum} number of sequential detections within a given time to { merger}, { assuming} that SLI would catch all the sources { that were previously observed by PTA-SKA}. { We then} consider the impact of the SLI low cut-off frequency, and of the spin of the remnants. Then for NGO, we compute the SNR at different values of the redshift $z$, and we finally adapt our previous estimates for sequential detection.  Section 4 sums up the conclusions; the appendix attempts to sketch some of the operational scenarios that may lie ahead.  We refer to the total mass of the binary $M$, normalised to $1 + z$, where $z$ is the { red shift} (Hughes 2002). The mass enters in the orbit evolution equation with the { time scale} $Gm/c^3$, the { time scale} being red shifted. The consequence\\footnote{Petiteau et al. (2011),  discuss the feasibility of breaking the degeneracy with electromagnetic counterparts and propose enforcing statistical consistency.} is that a binary at $z=1.5$ and of mass $2\\cdot 10^8~M_\\odot$ is equivalent to a binary at $z = 0.2$ and of mass $4.17\\cdot 10^8~M_\\odot$. Herein, given the range of $z$, the factor $(1+z)$ introduces an uncertainty of less than $2.25$ on the mass determination.  We use the terms of observation and detection when referring to PTA-SKA and SLI, respectively, { while} sequential detection implies both observation by PTA-SKA and detection by SLI. Finally, the term `event' refers mostly to astrophysical phenomena, or it is used whenever a specific labelling { is not intended}.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We have studied SMBHBs { with a mass range} of $2 \\cdot 10^8 - 2 \\cdot 10^9~M_\\odot$. The masses are normalised to a $(1+z)$ factor, the { red shift} $z$ lying between $z=0.2$ and $z=1.5$), where the individually detectable sources by PTA-SKA are expected to lie. For { a high PTA-SKA frequency cut-off of} $f_{\\mathit p}= 8 \\cdot 10^{-7}$ Hz, SMBHBs may pass from the pulsar timing to laser interferometry band in a period ranging from two months to eight years ($f_{\\mathit p}= 4 \\cdot 10^{-7}$ Hz, from thirteen months to fifty years). Furthermore, the source signals may be strong enough to be received by both { radio astronomical} and laser interferometric means.  { Unfortunately,} the astrophysical estimates act as the show-stoppers to sequential detection. Even a noiseless, extremely low frequency laser space interferometer, coupled to an highly performing pulsar timing of $1$ ns, will not allow to go beyond $4.7 ~\\cdot10^{-4}$ y$^{-2}$, or $3.3 ~\\cdot10^{-6}$ y$^{-2}$ sequential detections per year to { merger}  and per year of survey, for { the} most optimistic and pessimistic astrophysical { models}, respectively. We { are dealing with a} handful of individually detectable SMBHBs, and therefore it is no surprise that only a very tiny part of those sources merge within a time limit. Other factors { reduce the chances even further}: higher values of the SLI low cut-off frequency, low precision timing residuals, slowly rotating remnants, and { obviously SLI} noise. We have also found that the spread between the different astrophysical predictions { decreases} as timing precision improves, and  the SLI low cut-off frequency decreases.  Given the paucity of rates, major changes may come from radically new astrophysical models, or from the achievement of sub-nanosecond precision (predictions may then converge to a rate of $n \\cdot 10^{-1} y^{-2}$, per year to { merger}  and per year of survey). { For the models, the recent estimates on merger rates and gravitational waves amplitude by McWilliams et al. (2012) are encouraging.}  Complementarity between PTA-SKA and SLI should not be dismissed { out of hand}. Approaching the two bands would be helpful, and efforts to increment the availability of observation time may create favourable conditions to enlarge the PTA bandwidth. SLI should { remain interesting} for low frequencies in light of the proposed  higher frequency interferometers BBO (Phinney et al. 2003) and DECIGO (Seto et al. 2001). An SLI configuration having a cut-off frequency at $10^{-6}$ Hz was proposed by Bender (2004). %\\footnote{The sensitivity curve is given by $3 \\cdot 10^{-16}$ m/s$^2/\\sqrt{Hz}$ down to $0.1$ mHz, then going as $f^{-0.5}$ down to $0.01$ mHz, and finally as $f^{-1}$ from $0.01$ to $0.001$ mHz. As secondary note, the transit time slightly decreases due to the lower  $f_{\\mathit s}$ cut-off frequency.}.  Opportunities may also lie in inverse search. Ring-down signals observed by SLI may trigger specific searches by { PTA-SKA}, and hopefully allow to dig out the SMBHB from the background in the { PTA-SKA} band. { Another} opportunity may be provided by observation of PTA-SKA and the consequent evaluation of the SMBHB parameters (mass, spin), leading to a prediction when the merger would occur.  { In the absence in the immediate future of} a detector covering the gap { between PTA-SKA and SLI frequency bands}, these efforts { might be rewarded}.  Indeed, { the} detection of even a single event would be of paramount importance for the understanding of SMBHBs and of the astrophysical processes connected to their formation and evolution."
    },
    "1107/1107.0883_arXiv.txt": {
        "abstract": "{} { Only for very few $\\beta$~Cephei stars has the behaviour of the magnetic field been studied over the rotation cycle. During the past two years we have obtained multi-epoch polarimetric spectra of the $\\beta$~Cephei star V1449\\,Aql with SOFIN at the Nordic Optical Telescope to search for a rotation period and to constrain the geometry of the magnetic field. } { The mean longitudinal magnetic field is measured at 13 different epochs. The new measurements, together with the previous FORS\\,1 measurements, have been used  for the frequency analysis and the characterization of the magnetic field. } { V1449\\,Aql so far possesses the strongest longitudinal magnetic field of up to 700\\,G among the $\\beta$\\,Cephei stars. The resulting periodogram displays three dominant peaks with the highest peak at $f=0.0720$\\,d$^{-1}$ corresponding to a period $P=13\\fd893$. The magnetic field geometry can likely be described by a centred dipole with a polar magnetic field strength $B_{\\rm d}$ around 3\\,kG and an inclination angle $\\beta$ of the magnetic axis to the rotation axis of 76$\\pm$4$^{\\circ}$. As of today, the strongest longitudinal magnetic fields are detected in the $\\beta$\\,Cephei stars V1449\\,Aql and $\\xi^1$\\,CMa with large radial velocity amplitudes. Their peak-to-peak amplitudes reach $\\sim$90\\,km\\,s$^{-1}$ and $\\sim$33\\,km\\,s$^{-1}$, respectively. Concluding, we briefly discuss the position of the currently known eight magnetic $\\beta$\\,Cephei and candidate $\\beta$\\,Cephei stars in the Hertzsprung-Russell (H-R) diagram. } {} ",
        "introduction": "Before the CoRoT mission, the B1.5~II-III star V1449\\,Aql (=HD\\,180642) was known as a $\\beta\\,$Cephei pulsator with a high-amplitude radial mode of frequency 5.487\\,d$^{-1}$ (Waelkens et al.\\ \\cite{Waelkens1998}). White light photometry from space revealed the rich frequency spectrum of this target (Degroote et al.\\ \\cite{Degroote2009}) with the first detection of solar-like oscillations in a massive star (Belkacem et al.\\ \\cite{Belkacem2009}). Additional ground-based, multi-colour photometry and high-resolution spectroscopy were used to put constraints on the atmospheric parameters and the chemical abundances, and to identify three modes of low amplitudes (Briquet et al.\\ \\cite{Briquet2009}). Notably, the dominant mode does not behave sinusoidally, and the first velocity moment $\\langle v^1\\rangle$ has a peak-to-peak amplitude of about 90 km/s. This is the third largest peak-to-peak radial--velocity amplitude measured among  $\\beta~$Cephei stars, after BW\\,Vul (e.g.\\ Crowe \\& Gillet \\cite{CroweGillet1989}) and $\\sigma$~Sco (Mathias et al.\\ \\cite{Mathias1991}). A detailed NLTE abundance analysis showed a nitrogen excess of $\\sim$0.3\\,dex with respect to the average nitrogen abundance in B-type stars in the solar neighbourhood. Previous abundance studies of hot stars indicate that the appearance of such an excess may be linked to the presence of a magnetic field (e.g., Morel et al.\\ \\cite{Morel2006,Morel2008}; Przybilla \\& Nieva \\ \\cite{PrzybillaNieva2010}).    The first spectropolarimetric observations of V1449\\,Aql were obtained  with the multi-mode instrument FORS\\,1  at the VLT on two consecutive nights in 2007 and revealed a change in polarity from one night to the next and the presence of a weak magnetic field on the second night of observations (Hubrig et al.\\ \\cite{Hubrig2009}). In the present work we discuss 13 new spectropolarimetric observations acquired during the last two years at the 2.56\\,m Nordic Optical Telescope using the SOFIN echelle spectrograph. The new longitudinal magnetic field measurements are used to determine the stellar rotation period and to put constraints on the magnetic field geometry. ",
        "conclusions": "Using FORS\\,1/2 and SOFIN longitudinal magnetic field measurements collected in our recent studies, we were able to detect a weak mean longitudinal magnetic field of a few hundred Gauss in four $\\beta$\\,Cephei stars, $\\delta$\\,Cet, $\\xi^1$\\,CMa, 15\\,CMa, and V1449\\,Aql, and in two candidate $\\beta$\\,Cephei stars, $\\alpha$~Pyx and $\\epsilon$~Lup (Hubrig et al.\\ \\cite{Hubrig2006,Hubrig2009,Hubrig2011}). Before we started our systematic search for magnetic fields in pulsating B-type stars, a weak magnetic field was detected in two other $\\beta$\\,Cephei stars, in the prototype of the class, $\\beta$\\,Cep itself, by Henrichs et al.\\ (\\cite{Henrichs2000}) and in V2052\\,Oph by Neiner et al.\\ (\\cite{Neiner2003}). The strongest longitudinal magnetic fields were detected in the two $\\beta$\\,Cephei stars with the large radial velocity amplitudes, namely in V1449\\,Aql and $\\xi^1$\\,CMa, for which peak-to-peak amplitudes reach $\\sim$90\\,km\\,s$^{-1}$ and $\\sim$33\\,km\\,s$^{-1}$, respectively. The rather strong magnetic field in the atmosphere of these two stars and the rather low  $v$\\,sin\\,$i$ values, certainly make these stars the most suitable targets for studying  the impact of magnetic fields on stellar rotation, pulsations, element diffusion, and frequency patterns.  In the frequency analysis of the CoRoT data obtained for V1449\\,Aql (Degroote et al.\\ \\cite{Degroote2009}), the frequency $f=0.068790\\pm0.000393$\\,d$^{-1}$ is, within the errors, close to the frequency $f=0.069162$\\,d$^{-1}$ detected in our study of the periodicity using magnetic field measurements. Additional magnetic field measurements will be worthwhile for proving whether this photometric frequency corresponds to the rotation/magnetic period of this star.  \\begin{figure} \\centering \\includegraphics[height=0.35\\textwidth,angle=270]{17261f6a.eps} \\includegraphics[height=0.35\\textwidth,angle=270]{17261f6b.eps} \\caption{The position of the $\\beta$\\,Cephei and candidate $\\beta$\\,Cephei stars in the H-R diagram. Upper panel: The boundaries of the theoretical instability strips are calculated using the OP opacities. Lower panel: Boundaries are calculated using OPAL opacities. Full lines correspond to strips for metallicity $Z=0.01$ and dotted lines to strips with metallicity $Z=0.02$. The thick and thin lines correspond to the boundaries of the $\\beta$~Cephei and SPB instability regions, respectively. } \\label{fig:m6149} \\end{figure}   Our studies indicate that dipole models provide a satisfactory fit to the magnetic data and among the  presently known magnetic $\\beta$\\,Cephei stars, $\\xi^1$\\,CMa and V1449\\,Aql, possess the largest magnetic fields, with a dipole strength of several kG. The position of the currently known eight magnetic $\\beta$\\,Cephei and candidate $\\beta$\\,Cephei stars in the log\\,T$_{\\rm eff}$--log\\,$g$ diagram is presented in Fig.~\\ref{fig:m6149} together with the boundaries of the theoretical instability strips calculated for different metallicities ($Z=0.01$ and $Z=0.02$) and using the OP opacities (http://cdsweb.u-strasbg.fr/topbase/op.html, see also Miglio et al.\\ \\cite{Miglio2007}) and OPAL opacities (http://opalopacity.llnl.gov/opal.html). The stellar fundamental parameters and their literature sources are compiled in Table~\\ref{tab:fundamental}. We give in the same table the information on the dominant mode and the radial velocity amplitude, if known. The extent and the position of the instability strip are highly dependent on the opacities used and on the metallicity adopted. Recent studies indicate, however, a lower metallicity than $Z=0.02$ for early B-type stars in the solar neighbourhood with $Z=0.014$ (e.g.\\ Przybilla et al.\\ \\cite{Przybilla2008}). Although the sample of magnetic pulsating stars is still very small, and it is difficult to determine any trend in the distribution of the stars in the H-R diagram, it might seem that magnetic stars are clustered rather close to the instability strip boundaries (full lines in left panel of Fig.~\\ref{fig:m6149}). Certainly, additional future magnetic field studies of a larger sample of $\\beta$\\,Cephei stars are needed to study the indicated trend."
    },
    "1107/1107.0767_arXiv.txt": {
        "abstract": "{ Based on the turbulent convection model (TCM) of Li \\& Yang (2007), we have studied the characteristics of turbulent convection in the envelopes of $2$ and $5M_{\\odot}$ stars at the RGB and AGB phases. The TCM has been applied successfully in the whole convective envelopes including the convective unstable zone and the overshooting regions. We find that the convective motions become stronger and stronger when the stellar models are located upper and upper along the Hayashi line. In the convective unstable zone, we find that the turbulent correlations are proportional to functions of a common factor $(\\nabla-\\nabla_{ad})\\overline{T}$, which explains similar distributions of those correlations. For the TCM we find that if the obtained stellar structure of temperature is close to that of the MLT, the convective motion will be at a much larger velocity and thus more violent. However, if the turbulent velocity is adjusted to close to that of the MLT, the superadiabatic convection zone is much more extended inward and a lower effective temperature of the stellar model will be obtained. For the overshooting distance we find that the e-folding lengths of the turbulent kinetic energy $k$ in both the top and bottom overshooting regions decrease as the stellar model is located up along the Hayashi line, but both the extents of the decrease are not obvious. And the overshooting distances of the turbulent correlation $\\overline{u_{r}^{'}T^{'}}$ are almost the same for the different stellar models with a same set of TCM's parameters. For the decay way of the  kinetic energy $k$ we find that they are very similar for the different stellar models based on a same set of TCM's parameters, and there is a nearly linear relationship between $\\lg k$ and $\\ln P$ for the different stellar models. When $C_{s}$ or $\\alpha$ increases while the other parameters are fixed, the obtained linearly decaying distance will become longer. ",
        "introduction": "% \\label{sect:intro} Red giant branch (RGB) and asymptotic giant branch (AGB) stars are characterized by the convective motion in their envelopes, which can extend over an enormous range in mass (or radius). Because of huge scale of convection zones and small viscosity of the stellar material, turbulent flow instead of laminar one always occurs in their outer envelopes. Turbulent convection may result in many important effects in stars: mixing different elements to be homogeneous in the convection zone and adding fresh fuel to the nuclear burning region, dredging the internal material processed by H or He burning up to the stellar surface, and being as an important means for heat transfer. These effects significantly influence the structure and evolution of stars. Many observational phenomena appearing in the RGB and AGB stars are ascribed to the convection. For the RGB stars as an example, many of them show chemical anomalies (Pilachowski et al. 1993; Charbonnel 1995; Gratton et al. 2000; Kraft et al. 1993; Ram\\'{\\i}ez \\& Cohen 2002; Gratton et al. 2004; Busso et al. 2007; Recio-Blanco \\& de Laverny 2007). And we are still perplexed to the so-called carbon star mystery (Iben 1981; Iben 1975, 1977; Sackmann 1980; Lattanzio 1989; Hollowell \\& Iben 1988; Straniero et al. 1997; Herwig et al. 1997; Herwig 2005) for the AGB stars. Furthermore, the observed mass loss occurring in the RGB and AGB stars may be due to the turbulent pressure (Jiang \\& Huang 1997). To solve these problems we need to accurately deal with the convection and to learn details of its characteristics. \\parskip=.0in  For some RGB and AGB stars the convective motion near the surface of the stellar envelope are superadiabtic because of less efficient convective energy transport, and may sometimes become supersonic based on the mixing-length theory (MLT) (Deng \\& Xiong 2001). The MLT (B\\\"{o}hm-Vitense 1953, 1958) is commonly used to deal with the convection. However, due to its simplicity many shortcomings are found (Renzini 1987; Baker \\& Kuhfuss 1987; Spruit et al. 1990; Pederson et al. 1990), which thus lead to some confused and debatable problems. For example, convective overshooting (e.g. Saslaw \\& Schwarzschild 1965), semiconvection (e.g. Castellani et al. 1971; Sreenivasan \\& Wilson 1978) and the breathing convection (e.g. Castellani et al. 1985).  Therefore new convection models have been proposed successively, for example, non-local mixing length theories (Ulrich 1970; Maeder 1975; Bressan et al. 1981) and full-spectrum convection theory (Canuto \\& Mazzitelli 1991). However, they are still based on the framework of the MLT. A better theory to overcome the MLT's drawbacks is the turbulent convection models (TCMs), which are deduced directly from the Navier-Stokes equations, and can describe many turbulence properties (Xiong 1979, 1981, 1985, 1990; Canuto 1992, 1994, 1997; Xiong, Cheng \\& Deng 1997; Canuto \\& Dubovikov 1998; Canuto et al. 1996). However, the dynamic equations for turbulent correlations are expressed by higher order correspondents due to the inherent non-linearity of the Navier-Stokes equations. Therefore these equations are needed to be cut-off with some closure approximations in order to be applicable in calculations of the stellar structure and evolution. Nevertheless, some free parameters have to be introduced by the closure assumptions. Different closure assumptions correspond to different TCMs, and many of them may not be easy to be incorporated into a stellar evolution code. Recently a simple TCM proposed by Li \\& Yang (2007) was successfully applied to the solar models (Yang \\& Li 2007), in which the introduced parameters of the TCM are directly related to the corresponding turbulence physics. Some significant changes in the structure of the solar convection zone and better results compared to the solar p-mode observations are obtained. Shortly afterwards, Zhang \\& Li (2009) applied it successfully to the overshooting region of the solar convection zone. \\parskip=.0in  Up to now, the TCMs have seldom been applied to very large convective envelopes of stars, e.g. in the RGB or AGB stars, which are rather different from the solar environment. Furthermore, in view of the key roles of convection played in these stars, we choose three sites of stellar evolution to test the TCM of Li \\& Yang (2007). We try to find out convection characteristics for two stars of 5 and $2M_{\\odot}$ at the RGB and AGB phases, and to find out their dependence on the TCM's parameter. To analyze the properties of the turbulent correlations among velocity and temperature fluctuations, we separate the whole convection envelope (derived self-consistently from the TCM) into the convective unstable zone and the convective overshooting regions, respectively. We firstly describe the basic equations of the TCM in Section 2. The information of the evolutionary code and input physics are described in Section 3. In Section 4, the convection characteristics in the convective unstable zone and overshooting regions are presented and discussed, respectively. The dependence of the structure of the overshooting regions on the TCM's parameters is given in Section 5. Finally, some concluding remarks are summarized in Section 6. ",
        "conclusions": "Based on the TCM of Li \\& Yang (2007) we have obtained the characteristics of the turbulent convection in the RGB and AGB stars with huge convection zones in their envelopes. Some sets of suitable values of the TCM's parameters that have been used in the solar convection zone are adopted to make an exploratory application in two stars of $2$ and $5M_{\\odot}$ at the RGB and AGB phases. In practical calculations the TCM has been applied to the whole convective envelope including the overshooting regions. When analyzing the characteristics of the turbulent correlations, we separate the whole convective envelope into the convective unstable zone and the convective overshooting region. In the convective unstable zone, we find an approximation approach to associate the turbulent correlations explicitly, and can be used to explain the turbulent properties. Later, we make a comparison between the results of the TCM and those of the MLT on the velocity profile of turbulence and in influence on stellar structure and evolution. These obtained features of convective motion by the TCM are highly sensitive to the TCM's parameters. The main conclusions are summarized as follows:  \\hangafter=1\\setlength{\\hangindent}{2.8em}1) The non-local effect of the turbulent convection is only exhibited significantly around the boundaries of the convective unstable zone, especially around the upper boundary, which is in agreement with the results of Li \\& Yang (2007). The three correlations $\\sqrt{k}$, $\\overline{u_{r}^{'}T^{'}}$, and $\\overline{{T'}^{2}}$ are nearly peaked at the same place in the convective unstable zone for a same set of the TCM's parameters. It is because that all of them are directly correlated to a function of $(\\nabla-\\nabla_{ad})\\overline{T}$.   \\hangafter=1\\setlength{\\hangindent}{2.8em}2) For a fixed heat flux (closely related by $\\overline{u_{r}^{'}T^{'}}$) transferred by the convective motion, compared to the results of the MLT, the TCM will result in a larger turbulent velocity and thus more violent convective motion. Furthermore, for a nearly same profile of the convection velocity obtained by adjusting the value of the mixing length parameter $\\alpha$ for both the TCM and MLT, the superadiabatic convection zone will be extended much more inward for the TCM than for the case of the MLT and the stellar models based on the TCM will be of lower effective temperature.  \\hangafter=1\\setlength{\\hangindent}{2.8em}3) For the stellar models located up along the Hayashi line, we find that the convective motions become stronger and stronger. However, the e-folding lengths of $\\sqrt{k}$ in both the top and bottom overshooting regions decrease as the stellar model is located up along the Hayashi line. The overshooting length of $\\overline{u_{r}^{'}T^{'}}$ is found to be shorter than other turbulent quantities for a same set of the TCM's parameters in both overshooting regions, which is consistent with the results of Deng \\& Xiong (2008) and Zhang \\& Li (2009). And the overshooting distances of $\\overline{u_{r}^{'}T^{'}}$ in the bottom overshooting region are almost the same for different stellar models with a same set of the TCM's parameters. Therefore, the bottom convective overshooting has a similar effect on the stellar structure for all of the considered stellar models.  \\hangafter=1\\setlength{\\hangindent}{2.8em}4) The diffusive parameter $C_{s}$ plays an important role in the development of the isotropic turbulence in the top overshooting region, and the longer the isotropic turbulence will be obtained for larger  $C_{s}$. In the bottom overshooting region, the convection becomes more and more horizontally dominated for larger value of $C_{s}$, and the convection shows a much more complicated behavior when $C_{s}$ is small enough.  \\hangafter=1\\setlength{\\hangindent}{2.8em}5) The decaying ways of the turbulent kinetic energy $k$ are very similar for the different stellar models based on a same set of TCM's parameters. We find that there is a nearly linear relation between $\\lg k$ and $\\ln P$ in most of the overshooting regions. The slope of the decaying law seems only a function of the TCM's parameters. The values of $C_{s}$ and $\\alpha$ have a similar effect on the decaying distance of the turbulent kinetic energy, and the larger the values of them the slope of the decaying law becomes smaller and therefore the linearly decaying distance becomes longer.      \\normalem"
    },
    "1107/1107.5374_arXiv.txt": {
        "abstract": "Continuing work initiated in an earlier publication (Abe, ApJ, {\\bf 725} (2010) 787), we study the gravitational microlensing effects of the Ellis wormhole in the weak-field limit. First, we find a suitable coordinate transformation, such that the lens equation and analytic expressions of the lensed image positions can become much simpler than the previous ones. Second, we prove that two images always appear for the weak-field lens by the Ellis wormhole. By using these analytic results, we discuss astrometric image centroid displacements due to gravitational microlensing by the Ellis wormhole. The astrometric image centroid trajectory by the Ellis wormhole is different from the standard one by a spherical lensing object that is expressed by the Schwarzschild metric. The anomalous shift of the image centroid by the Ellis wormhole lens is smaller than that by the Schwarzschild lens, provided that the impact parameter and the Einstein ring radius are the same. Therefore, the lensed image centroid by the Ellis wormhole moves slower. Such a difference, though it is very small, will be in principle applicable for detecting or constraining the Ellis wormhole by using future high-precision astrometry observations. In particular, the image centroid position gives us an additional information, so that the parameter degeneracy existing in photometric microlensing can be partially broken. The anomalous shift reaches the order of a few micro arcsec. if our galaxy hosts a wormhole with throat radius larger than $10^5$ km. When the source moves tangentially to the Einstein ring for instance, the maximum position shift of the image centroid by the Ellis wormhole is $0.18$ normalized by the Einstein ring radius. For the same source trajectory, the maximum difference between the centroid displacement by the Ellis wormhole lens and that by the Schwarzschild one with the same Einstein ring radius is $-0.16$ in the units of the Einstein radius, where the negative means that the astrometric displacement by the Ellis wormhole lens is smaller than that by the Schwarzschild one. ",
        "introduction": "A peculiar feature of general relativity is that the theory admits a nontrivial topology of a spacetime. A solution of the Einstein equation that connects distant points of space--time was introduced by \\citet{ein35}. This solution called the {\\it Einstein--Rosen bridge} was the first solution to later be referred to as a wormhole. Initially, this type of solution was just a trivial or teaching example of mathematical physics. However, \\citet{morr88} proved that some wormholes are \"traversable\"; i.e., space and time travel can be achieved by passing through the wormholes. They also showed that the existence of a wormhole requires exotic matter that violates the null energy condition. The existence of wormholes, though they are very exotic, has not been ruled out in theory. Inspired by the Morris--Thorne paper, a number of theoretical works (see \\citet{vis95, lob09} and references therein) on wormholes have been done. The curious nature of wormholes such as time travel, negative energy conditions, space--time foams, and growth of a wormhole in an accelerating universe have been studied. In spite of enthusiastic theoretical investigations, studies searching for observational evidence of the existence of wormholes are scarce. Only a few attempts have been made to show the existence or nonexistence of wormholes in our universe \\citep{cram95, tor98, saf01, bog08}.  A possible observational method that has been proposed to detect or exclude the existence of wormholes is the application of optical gravitational lensing, since the light ray propagation is sensitive to a local spacetime geometry. The gravitational lensing of wormholes was pioneered by \\citet{cram95}, who inferred that some wormholes show \"negative mass\" lensing. They showed that the light curve of the negative-mass lensing event of a distant star has singular double peaks. Several authors subsequently conducted theoretical studies on detectability \\citep{saf01, bog08}. Another gravitational lensing method employing gamma rays was proposed by \\citet{tor98}, who postulated that the singular negative-mass lensing of distant active galactic nuclei causes a sharp spike of gamma rays and may be observed as double-peaked gamma-ray bursts. They analyzed BATSE data to put an upper limit to the negative mass density $O(10^{-36}) \\:\\mbox{g cm}^{-3}$ in the form of wormholelike objects.  There have been several recent works \\citep{sh04, per04, nan06, rah07, dey08, abe10, asada11} on the gravitational lensing of wormholes as structures of space--time. Such studies are expected to unveil lensing properties directly from the space--time structure. One study \\citet{dey08} calculated the deflection angle of light due to the Ellis wormhole, whose asymptotic mass at infinity is zero. The massless wormhole is particularly interesting because it is expected to have unique gravitational lensing effects. This type of wormhole was first introduced by \\cite{ell73} as a massless scalar field. Later, \\citet{morr88} studied this wormhole and proved it to be traversable. The dynamical feature was studied by \\cite{shi02}, who showed that Gaussian perturbation causes either explode to an inflationary universe or collapse to a black hole. \\cite{das05} showed that the tachyon condensate can be a source for the Ellis geometry. \\citet{abe10} provided a method to calculate light curves of the gravitational microlensing of the Ellis wormhole in the weak-field limit. This result has been discussed as one example of corrections to the standard formula of the deflection angle by \\cite{asada11}.  The main results of this paper are: (1) We derive analytic expressions for calculations of the wormhole lensing. (2) We show the astrometric image centroid trajectory by the Ellis wormhole lens. Studies of centroid displacements of lensed images have been limited within the Schwarzschild lens \\citep{Walker,MY,HOF,SDR,JHP,asada02,HL}. In Section 2, we discuss gravitational lensing by the Ellis wormhole in the weak-field limit. We use a suitable coordinate transformation, such that the lens equation to determine image positions can take a much simpler form than the previous one derived by \\citet{abe10}. By using this method, calculations of the gravitational microlensing by the Ellis wormhole are shortened. In addition, we prove that two images always appear for the weak-field lens by the Ellis wormhole. In section 3, we discuss astrometric image centroid displacements due to gravitational microlensing by the Ellis wormhole. The results are summarized in Section 4. ",
        "conclusions": "We studied the gravitational microlensing effects of the Ellis wormhole in the weak-field limit. First, we performed a suitable coordinate transformation, such that the lens equation and analytic expressions of the lensed image positions can become much simpler than the previous ones. Second, we proved that two images always appear for the weak-field lens by the Ellis wormhole. By using these analytic results, we investigated astrometric image centroid displacements due to gravitational microlensing by the Ellis wormhole. The anomalous shift of the image centroid by the Ellis wormhole lens is smaller than that by the Schwarzschild lens, provided that the impact parameter and the Einstein ring radius are the same. Therefore, the lensed image centroid by the Ellis wormhole moves slower.  Studies of astrometric image centroid displacements due to another type of wormholes or nontrivial topology of spacetimes are left as a future work."
    },
    "1107/1107.3528_arXiv.txt": {
        "abstract": "Space-based gravitational-wave (GW) detectors, such as LISA or a similar ESA-led mission, will offer unique opportunities to test general relativity. We study the bounds that space-based detectors could realistically place on the graviton Compton wavelength $\\lambda_g=h/(m_g c)$ by observing multiple inspiralling black hole (BH) binaries. We show that while observations of {\\it individual} inspirals will yield mean bounds $\\lambda_g\\sim 3\\times 10^{15}$~km, the {\\it combined} bound from observing $\\sim 50$ events in a two-year mission is about ten times better: $\\lambda_g\\simeq 3\\times 10^{16}$~km ($m_g\\simeq 4\\times 10^{-26}$~eV). The bound improves faster than the square root of the number of observed events, because typically a few sources provide constraints as much as three times better than the mean. This result is only mildly dependent on details of BH formation and detector characteristics. The bound achievable in practice should be one order of magnitude better than this figure (and hence almost competitive with the static, model-dependent bounds from gravitational effects on cosmological scales), because our calculations ignore the merger/ringdown portion of the waveform. The observation that an ensemble of events can sensibly improve the bounds that individual binaries set on $\\lambda_g$ applies to any theory whose deviations from general relativity are parametrized by a set of global parameters. ",
        "introduction": "Tight dynamical bounds on the graviton mass can be obtained using GW observations in space. This is due to two reasons: (1) the larger mass of observable BH binaries, as compared to ground-based GW observations; (2) the statistical increase in the bound that would result from observing {\\it populations} of individually resolved binaries. To illustrate point (1), in this section we review existing work on massive graviton bounds from GW observations of individual binaries with LISA. In section \\ref{sec:pops} we will present the first attempt to quantify the statistical improvement of bounds on $\\lambda_g$ that could be achievable in reality by observing several events in the context of hierarchical models of massive BH formation.  Will \\cite{Will:1997bb} first pointed out that interesting bounds on $\\lambda_g$ could come from a careful monitoring of the phase of GWs emitted by binaries of compact objects, such as BHs and/or neutron stars.  In hypothetical massive graviton theories, the GW damping formulae and dispersion relation would be modified. As a consequence, the GW phasing $\\Psi_{\\rm GR}(f)$ would acquire an additional term: \\be \\Psi_{\\rm MG}(f)=\\Psi_{\\rm GR}(f)- \\beta_g (\\pi Mf)^{-1}\\,, \\ee where $\\beta_g\\equiv \\pi^2 DM/[(1+z)\\lambda_g^2]$ and $D$ is a distance parameter, similar to (but not quite the same as) the luminosity distance $D_L$ (here and below we assume a $\\Lambda$CDM model with $H_0=70$ km$\\cdot$s$^{-1}$Mpc$^{-1}$, $\\Omega_M=0.3$, $\\Omega_M=0.7$).  Using the restricted post-Newtonian approximation and neglecting binary spins, Will estimated that stellar-mass binary inspirals to be observed with LIGO would yield a bound $\\lambda_g\\simeq 5\\times 10^{12}$~km, only slightly better than Solar System bounds. For an ``optimal'' LISA system, i.e. an equal-mass BH binary of total mass $M=2\\times 10^6~M_\\odot$ at $D_L=3$~Gpc ($z\\simeq 0.5$), the bound would be four orders of magnitude better: $\\lambda_g\\simeq 5\\times 10^{16}~{\\rm km}\\simeq 1.6~{\\rm kpc}$.  These initial estimates were refined in various papers. Will and Yunes \\cite{Will:2004xi} showed that bounds from binary BH inspirals at $D_L=3$~Gpc would range between $10^{15}$~km and $5\\times 10^{16}$~km for $M$ in the range $10^4-10^7~M_\\odot$.  They found that the bound on $\\lambda_g$ is proportional to $\\sqrt{L}$ (where $L$ is the LISA armlength) and to (LISA acceleration noise)$^{-1/2}$, and that it scales in the following way: \\be \\lambda_g\\propto \\left(\\f{D}{(1+z)D_L}\\right)^{1/2} \\f{{\\cal M}^{11/12}}{S_0^{1/4}f_0^{1/3}}\\,, \\ee where ${\\cal M}=(m_1 m_2)^{3/5}M^{-1/5}$ is the ``chirp mass'', $S_0$ (in Hz$^{-1}$) sets the scale of the noise power spectral density (PSD) of the detector, and $f_0$ is a characteristic ``knee'' frequency where the PSD has a minimum. Distance dependence is weak because the effect of the massive graviton and measurement errors both grow with distance.  Berti {\\it et al.} \\cite{Berti:2004bd} performed a more detailed Monte Carlo calculation of LISA's parameter estimation capabilities. They still considered restricted post-Newtonian inspiral waveforms, but they used Cutler's model \\cite{Cutler:1997ta} to take into account the motion of the detector (all previous analyses assumed a {\\it static} LISA constellation). In this paper we will do the same. In Cutler's model, LISA can be seen as one (two) independent ``LIGO-like'' Michelson interferometers depending on whether 4 (5/6, respectively) laser links are available between the three satellites forming the constellation.  For the ``optimal binary'' (a nonspinning, equal-mass binary with $M=2\\times 10^6~M_\\odot$ at $D_L=3~$Gpc), Monte Carlo averages over the binary position and orientation yield a bound $\\lambda_g= 5.0\\times 10^{16}(3.7\\times 10^{16})$~km when using two (one) Michelsons, in excellent agreement with Will's original results. For masses in the range $M=2\\times 10^4 - 2\\times 10^7~M_\\odot$, the mean bound using two Michelsons at $D_L=3~$Gpc is in the range $\\sim 5\\times 10^{15}-7\\times 10^{16}$~km.  Later work used mostly Cutler's model and investigated the effect of more complex inspiral waveforms. Arun and Will \\cite{Arun:2009pq} studied the effect of higher harmonics and PN amplitude corrections. They found that bounds on $\\lambda_g$ improve by a factor of a few for systems of total mass $\\gtrsim 10^6~M_\\odot$, and that this improvement is more significant for binaries with large mass ratios. If one ignores spin precession, adding spins to the waveforms generally introduces degeneracies between the binary parameters, degrading parameter estimation accuracy \\cite{Berti:2004bd}. Stavridis and Will \\cite{Stavridis:2009mb} showed that modulations induced by spin precession can break these degeneracies. For the ``optimal binary'', for example, the average bound including spin precession is $\\lambda_g\\sim 5\\times 10^{16}$~km, basically the same as in the case of nonspinning binaries. Yagi and Tanaka \\cite{Yagi:2009zm} included {\\it both} spin precession and eccentricity. Performing simulations for a $(10^7+10^6)~M_\\odot$ binary at 3~Gpc, they found an average bound $\\lambda_g=3.1\\times 10^{16}$~km. This is again consistent with \\cite{Berti:2004bd} to within a factor 2.  More recently, Keppel and Ajith \\cite{Keppel:2010qu} revisited this problem using phenomenological waveforms for nonspinning BH binaries that include the merger/ringdown phase. Their analysis of LISA bounds is similar to the original work by Will \\cite{Will:1997bb} in that it ignores the motion of the detector, using an effective non-sky-averaged noise PSD. This simplification",
        "conclusions": "We assessed the capability of future space-based interferometers, such as ``Classic LISA'' and the proposed ESA-led ``New LISA'', to constrain the mass of the graviton by combining observations of a {\\em population} of massive BH binaries.  We found that: (1) by using a population of merging BH binaries we can obtain a bound on $\\lambda_g$ that is $\\sim 10$ times better than the mean bound on individual observations; (2) quite independently of the detector's design and of details of the massive BH formation models, the combined bound from inspiral observations will be $\\lambda_g\\simeq 3\\times 10^{16}~{\\rm km}$. This figure is likely to underestimate the bound achievable in practice by about one order of magnitude, as we have ignored the merger and ringdown portion of the waveform \\cite{Keppel:2010qu}, but further work is required to confirm this expectation.  In conclusion, space-based observations of a population of merging BHs should set bounds in the range $\\lambda_g\\in [2\\times 10^{16}\\,,10^{18}~{\\rm km}]$ on the graviton Compton wavelength, depending on details of the detector and on the specific waveform model used to set the bounds. This is comparable to the (static and model-dependent) bounds from cosmological-scale observations quoted in Table \\ref{bounds} but it is very different in nature, because gravitational radiation tests the dynamical regime of Einstein's general relativity.  \\noindent {\\bf \\em Acknowledgments.} E.B. is supported by NSF Grant No. PHY-0900735 and by NSF CAREER Grant No. PHY-1055103. J.G.'s work is supported by the Royal Society. The authors wish to thank Martin Elvis and the ``New LISA'' Science Performance Task Force for discussions, Marta Volonteri for sharing her BH formation models, and the Aspen Centre for Physics (where this work was started) for providing a very stimulating environment."
    },
    "1107/1107.0001_arXiv.txt": {
        "abstract": "Nuclear inflows of metal--poor interstellar gas triggered by galaxy interactions can account for the systematically lower central oxygen abundances observed in local interacting galaxies.  Here, we investigate the metallicity evolution of a large set of simulations of colliding galaxies.  Our models include cooling, star formation, feedback, and a new stochastic method for tracking the mass recycled back to the interstellar medium from stellar winds and supernovae.  We study the influence of merger--induced inflows, enrichment, gas consumption, and galactic winds in determining the nuclear metallicity.  The central metallicity is primarily a competition between the inflow of low--metallicity gas and enrichment from star formation.  An average depression in the nuclear metallicity of $\\sim0.07$ is found for gas--poor disk--disk interactions.  Gas--rich disk--disk interactions, on the other hand, typically have an enhancement in the central metallicity that is positively correlated with the gas content.  The simulations fare reasonably well when compared to the observed mass--metallicity and separation--metallicity relationships, but further study is warranted. ",
        "introduction": "The nuclear metallicities of star-forming galaxies are characterized by a mass-metallicity relation \\citep[hereafter, MZ][]{Lequeux1979, Rubin1984,Tremonti2004}.  Contributing to the scatter in the MZ relation are interacting galaxies, which are consistently lower in central metallicity than non--interacting systems of equivalent mass, as first noted by \\citet{KGB06}, and later confirmed for ultraluminous infrared galaxies (ULIRGs) \\citep{Rupke2008}, close pairs in SDSS \\citep{SloanClosePairs,MichelDansac2008,Peeples2009,SolA2010}, and local, low mass systems \\citep{EktaChengalur2010}.  Observations such as these are most naturally explained as the result of vigorous merger--induced inflows of gas \\citep{BH91,BH96} which rearrange the initial metallicity gradient \\citep{Shields1990,BelleyRoy1992,Z94,MA04} and ``dilute'' the nuclear metallicity \\citep{KGB06}.  In this sense, the same gas that drives nuclear starbursts \\citep{MH94a,MH96,Iono04} and triggers central AGN activity and black hole growth \\citep{DSH05,SDMHModel,HH06, Hopkins07}, also results in suppressed nuclear metallicity.  Simulations of merger-driven inflows of gas naturally predict a flattening of the initial metallicity gradient~\\citep{Perez2006,Rupke2010,PerezScoop}, which has now been observed in a number of colliding systems \\citep{KewleyGradients2010,RupkeGradients2010}.  However, more than hydrodynamics are at play in determining the nuclear metallicity evolution.  To further our understanding, we should consider ongoing star formation with associated chemical enrichment, feedback from star formation and AGN activity, and the interchange of material between the stellar and gaseous phases as stars are born and later return material to the interstellar medium.  The four main processes responsible for the evolution of the nuclear metallicity are gas inflows, chemical enrichment from star formation, gas consumption, and galactic outflows.  These effects compete with one another to influence the nuclear metallicity, making it difficult to determine their relative importance a priori.  Numerical simulations have only recently been used to quantify the detailed impact of these effects on metallicity gradients and nuclear metallicities of interacting pairs.  \\citet{Rupke2010} simulated colliding galaxies without chemical enrichment to explore dynamically induced changes in metallicity gradients, finding that a drastic flattening can occur, accompanied by a drop in the nuclear metallicity.  \\citet{DilutionPeak2010} and~\\citet{PerezScoop} performed simulations with star formation and chemical enrichment and found similar results to \\citet{Rupke2010}, but noted that the dip in the nuclear metallicity values can be counteracted by chemical enrichment from star formation.  These simulations have refined our knowledge of the depressed MZ relation.  In this paper, we explore the evolution of the nuclear metallicity during mergers using numerical simulations which include cooling, star formation, stellar feedback, and black hole growth and AGN feedback. Our approach allows us to systematically investigate the importance of gas inflows, chemical enrichment from star formation, gas consumption as a result of star formation, and galactic outflows.  In order to unambiguously determine the role of metal enrichment, we have developed a stochastic method to recycle stellar particles back to the interstellar medium without requiring inter-particle mass mixing.  The stochastic gas recycling method, which is designed by analogy to the widely used stochastic star formation method, has attributes that are distinct from kernel weighted mass return, making it particularly useful for our study.  We consider a range of situations with systematically varied parameters to enhance our understanding of metallicity evolution as a natural extension of the merger process.  We find that, for systems modeled after those in the local Universe, the main driver behind changes to the nuclear metallicity is the flow of low metallicity gas to the nuclear regions, or metallicity dilution.  However, we identify two previously under-appreciated effects that influence the strength of this dilution: locking of gas and metals in the stellar phase, and gas and metal removal via stellar-driven winds.  We also find that, for systems modeled after high redshift galaxies, the main driver of the nuclear metallicity shifts to chemical enrichment.  In these simulations, the nuclear metallicity increases, contrary to what is observed in the local Universe, suggesting that the interacting galaxy MZ relation may evolve differently at high redshifts.  We compare our simulations directly to observations by synthesizing a population of progenitor galaxies with properties consistent with observed samples and show that the empirical depression in the close-pair mass-metallicity relation can be reproduced while accounting for chemical enrichment.  We find good agreement between our simulations and observations for both the interacting galaxy mass-metallicity relation and separation-metallicity relation. ",
        "conclusions": "Using numerical simulations, we have investigated the impact that mergers have on galactic nuclear metallicity.  Our models include the capability of describing star formation, chemical enrichment, gas recycling, and star-- formation driven winds.  The analysis performed here relies upon the ability to accurately track the spatial and temporal return of metals to the ISM enabled by our stochastic gas--recycling algorithm within GADGET (see \\S~\\ref{sec:GasRecycling}).  One of the primary results of this work is to reinforce the notion that galaxy mergers, and their attending gravitational tidal forces, generate significant inflows of gas shown explicitly in Figures~\\ref{fig:ZTorientations} and \\ref{fig:ZTavg} ~\\citep[see also][]{Rupke2010, DilutionPeak2010, PerezScoop}. These inflows transport metal--poor gas into the nuclear region and reinforce the generic connection between close galaxy interactions and nuclear metallicity dilution -- no matter how many close passages occur.  Merger--induced gaseous inflows are, however, not the only factor influencing the nuclear metallicity, and this paper quantifies the competing effects of enrichment from star formation (\\S~\\ref{sec:DilutionAndEnrichment}), gas consumption (\\S~\\ref{sec:GasConsumption}), and galactic winds (\\S~\\ref{sec:GalacticWinds}).  Our models provide a physical template to understand nuclear metallicity evolution and lead us to conclude that merger--induced metallicity dilution and chemical enrichment from star formation drive the nuclear metallicity, while gas consumption and galactic winds play a secondary role to modulate the efficiency of the primary processes.  We have made a distinct effort in this work to directly compare our models to observations, allowing us to validate the physical evolutionary model that we have presented and associate observable trends with the various stages of the merger evolution.   We find that the merger--induce depressions of nuclear metallicity are typically $\\sim0.07$ dex, an offset similar to that observed to the mass--metallicity relation for close pairs in SDSS~\\citep{SloanClosePairs}. The models also demonstrate that central depressions of metallicity are time--dependent owing to the specific merger orbit, the galaxies' spin--orbit coupling, and varying structural properties of the galaxies.  When compared to the observed separation--metallicity relation of~\\citet{KGB06}, the models generally show a similar trend to produce lower central metallicities at smaller separations.  The precise role of interactions is difficult to discern, however, because of the significant scatter in the relations (see \\ref{fig:SZ} and the discussion in~ \\S\\ref{sec:SZ}).  Future comparisons in which galaxies are morphologically classified according to merger stage, or separated according to IR luminosity or star--formation rate may be better suited to isolated the true effects of the interaction.  This work also demonstrates that the progenitor gas content can have a profound influence on the resulting merger--induced metallicity dilution. Mergers between gas--rich progenitors can yield systematic metallicity enhancement, as opposed to metallicity dilution (see \\S~\\ref{sec:MergerGasFraction}, and specifically Fig.~\\ref{fig:GasFractionHist}).  In our picture, central metallicity enrichment comes from vigorous merger--induced star formation that can compensate for, and eventually overcome the merger--induced inflow of metal--poor gas.  This view is slightly different than put forth in recent work by \\citet{PerezScoop}, which argues that disk instabilities and clump formation drive rapid star formation and subsequent metal enrichment.  But once these processes occur, which is typically early in the merger process, the story becomes similar to when the galaxies have low gas content.  By studying the detailed correlations between merger stage, pair separation, central metallicity, and metallicity gradients, in both observations and the models, future work will allow for a better understanding of the (merger and enrichment) processes at work."
    },
    "1107/1107.0830_arXiv.txt": {
        "abstract": "{The X-ray dim isolated neutron stars (XDINSs) are peculiar pulsar-like objects, characterized by their very well Planck-like spectrum. In studying their spectral energy distributions, the optical/UV excess is a long standing problem. Recently, Kaplan et al. (2011) have measured the optical/UV excess for all seven sources, which is understandable in the resonant cyclotron scattering (RCS) model previously addressed. The RCS model calculations show that the RCS process can account for the observed optical/UV excess for most sources . The flat spectrum of RX J2143.0+0654 may due to contribution from bremsstrahlung emission of the electron system in addition to the RCS process. ",
        "introduction": "The seven X-ray dim isolated neutron stars (XDINSs) are puzzling pulsar-like compact stars. They are characterized by their very well Planck-like spectrum, and then they are good specimen for neutron star cooling, atmospheric and equation of state studies (Tong et al. 2007,2008; review Turolla 2009). Additionally, these sources may also have some relations with other classes of pulsar-like objects, e.g., magnetars and rotating radio transients, etc (Tong et al. 2010).  The spectral energy distributions of XDINSs show that their optical and UV emissions are above the extrapolation of the X-ray blackbodies, i.e., optical/UV excess. This excess of XDINSs is a long standing problem which may be related to the equation of state of neutron stars (Xu 2002, 2009). Previously, only two sources (RX J1856.5-3754 and RX J0720.4-3125) have both optical and UV data (van Kerkwijk \\& Kaplan 2007). Recently, Kaplan et al. (2011) have presented results of the optical and UV emission for all seven sources. They also find that the source RX J2143.0+0654 has a very flat spectrm ($F_{\\nu} \\propto \\nu^{\\alpha}$, $\\alpha \\sim 0.5$). This new observation requires updated modeling of existing theoretical models.  Tong et al. (2010) have considered the resonant cyclotron scattering (RCS) in pulsar magnetospheres as the origin of optical/UV excess in XDINSs, i.e., the RCS model. This model can explain the observations of RX J1856.5-3754 and RX J0720.4-3125 (Tong et al. 2010). Considering the recent observational progress, we apply the RCS model to all seven XDINSs in this paper.  Section 2 is summary of the RCS model. Its application to XDINSs is treated in Section 3. Discussions and conclusions are presented in section 4. ",
        "conclusions": "In the discussion section of Kaplan et al. (2011), they also mentioned the RCS process, which they refer to the paper of Lyutikov \\& Gavriil (2006). However, as pointed in Tong et al. (2010), the one dimensional treatment of Lyutikov \\& Gavriil (2006) can only result in up scattering. While the optical/UV excess of XDINSs is of ``soft excess'' in nature, which requires down scattering process. This is the reason that Tong et al. (2010) modeled the RCS process three dimensionally and employed the Kompaneets equation method.  The RCS model can explain the spectral energy distributions of XDINSs in a wide parameter space. Meanwhile, since the surface X-ray photons are scattered by a shell of electrons, this will result in a low pulsation amplitude naturally. The X-ray spectra are Planck-like with no high energy tails. This may imply that XDINSs are possibly quark stars (Xu 2002,2009). At the same time, the possible features may be proton cyclotron lines or electron cyclotron lines far in the magnetosphere (Turolla 2009).  From figure 3 in Kaplan et al. (2011) we see that for RX J2143.0+0654 we have only two data points at present. Future more optical and UV data will tell us whether its spectrum is really so flat or not. Moreover, with more data points we can know whether there is curvature in the optical/UV spectrum and whether there are more spectra components in addition to the power law component. The electron thermal bresstrahlung component has a flat spectrum in the IR/optical range, with a flux about $0.1\\mu \\mathrm{Jy}$ for RX J2143.0+0654. This flux is much lower than the $H$-band flux upper limit ($1.6\\mu \\mathrm{Jy}$ at $1.65 \\,\\mu m$, Lo Curto et al. 2007). Therefore in order to verify whether the flat spectrum of RX J2143.0+0654 is due to thermal bremsstrahlung process, deeper optical/UV and IR observations are required.  The electron system cools down via the bremsstrahlung process (possibly also via free-bound transitions, i.e., recombination process). Therefore, the central star has to continually heat the electron system and ionize it. This is similar to the ``Str$\\ddot{\\mathrm{o}}$mgren sphere'' around early-type stars (Dyson \\& Williams 1980). The size of the ionized region is $(3S/4\\pi N_e^2 \\beta_2)^{1/3}$ (Dyson \\& Williamas 1980), where $S$ is the ionizing photon number flux of the central star, $N_e$ is the electron number density (assuming fully ionized hydrogen plasma), $\\beta_2$ is the recombination coefficient. Similar calculations show that the size of the ionized region is $\\sim 10^{10}\\,\\mathrm{cm}$, to order of magnitude consistent with the outer radius used in table 1."
    },
    "1107/1107.2142_arXiv.txt": {
        "abstract": "The radial distributions of temperature, density, and gas entropy among cool-core clusters tend to be quite similar, suggesting that they have entered a quasi-steady state. If that state is regulated by a combination of thermal conduction and feedback from a central AGN, then the characteristics of those radial profiles ought to contain information about the spatial distribution of AGN heat input and the relative importance of thermal conduction.  This paper addresses those topics by deriving steady-state solutions for clusters in which radiative cooling, electron thermal conduction, and thermal feedback fueled by accretion are all present, with the aim of interpreting the configurations of cool-core clusters in terms of steady-state models.  It finds that the core configurations of many cool-core clusters have entropy levels just below those of conductively balanced solutions in which magnetic fields have suppressed electron thermal conduction to $\\sim 1/3$ of the full Spitzer value, suggesting that AGN feedback is triggered when conduction can no longer compensate for radiative cooling.  And even when feedback is necessary to heat the central $\\sim 30$~kpc, conduction may still be the most important heating mechanism within a cluster's central $\\sim 100$~kpc. ",
        "introduction": "\\setcounter{footnote}{0}  The cores of galaxy clusters have lessons to teach about accretion of hot intergalactic gas onto massive galaxies and the feedback processes that limit star formation within those galaxies.    Many of the most massive galaxies in the universe can be found at the centers of galaxy clusters.  A cluster's central galaxy is often the brightest cluster galaxy (BCG), and the majority of central galaxies exhibit little or no star formation.  Minimal star formation is understandable for central galaxies in which the hot intracluster gas cannot radiate its thermal energy within a Hubble time but is harder to understand in clusters with shorter central cooling times.  Star formation is generally seen in the centers of galaxy clusters only when the central cooling time is $\\lesssim 1$~Gyr \\citep{mo92,Cardiel+98,Rafferty+08,Bildfell+08}.  And even then, the star formation rate is usually $\\lesssim 10$\\% of the rate one would naively infer from the apparent cooling rate of the intracluster medium \\citep{ODea+08}.  Feedback from a central active galactic nucleus (AGN) is the favored mechanism for suppressing cooling and star formation at the centers of these clusters \\citep[see][and references therein]{mn07}.  Relativistic plasma flowing from the AGN apparently evacuates large cavities within these clusters and in at least some cases drives shocks into the surrounding X-ray emitting medium \\citep[e.g.][]{Forman+07,Randall+11}.  The energy introduced is comparable to that radiated from the ambient hot gas, making the hypothesis of AGN feedback seem quite plausible \\citep[e.g.][]{Birzan+04}.  Yet, many questions about the AGN feedback mechanism remain unanswered.  Two of the most pressing are:  (1) how is the energy output of the AGN so well tuned to match the cooling rate of the surrounding hot gas, and (2) how does the AGN manage to distribute its energy throughout the core so that heating nearly balances cooling everywhere?  These questions arise because the AGN accretion rate would seem to depend on conditions in a neighborhood only a few tens of parsecs in size, far smaller than the $\\sim 100$~kpc cluster core that must be heated.  Furthermore, the entropy gradients of cluster cores with short central cooling times tend to increase gradually from $10-100$~kpc.  A large AGN outburst that dumped most of its energy within a $\\sim 10$~kpc region would overheat the center of the cluster, producing an entropy inversion \\citep[e.g.][]{vd05}.  Instead, the energy source that offsets cooling somehow balances the cooling rate out to $\\sim 100$~kpc without disrupting the gradually increasing radial entropy gradient.  Some implementations of AGN feedback in galaxy-cluster simulations appear promising \\citep[e.g.,][]{bs09,Dubois+10,Fabjan+10,McCarthy+11}, but the problem is far from being solved.  A closer look at the data suggests that clusters with central cooling times $\\lesssim 1$~Gyr do indeed achieve a quasi-steady balance between heating and cooling that sometimes allows star formation to proceed at moderate rates of $1-10 M_\\odot \\, {\\rm yr}^{-1}$.  Gas temperatures at $\\sim 10$~kpc radii in clusters with short central cooling times are 2--3 times cooler than those at  $\\sim 100$~kpc, which has led such clusters to be called ``cool-core\" clusters.  Gas densities decline by a factor $\\sim 10$ over the same range in radius.  And even though gas density and temperature profiles in these systems differ in normalization, their profiles of entropy ($K = P \\rho^{-5/3} \\propto kTn_e^{-2/3}$) as a function of radius are all quite similar \\citep{Donahue+06,Cavagnolo+09,sop09}.  A cluster's entropy distribution is of interest because it determines the temperature and density profiles of a cluster in hydrostatic equilibrium (see \\citet{Voit+02} and \\S~2).  Mergers and accretion shocks driven by gravitational structure formation produce a family of clusters with self-similar entropy profiles, and deviations from those gravitationally generated profiles reflect the action of non-gravitational cooling and heating processes \\citep[e.g.][]{Voit+03}.  The purpose of this paper is to explore why one particular set of density, temperature, and entropy profiles is favored among cool-core clusters.  Its mode of exploration is to outline the properties of various families of steady-state solutions for hot gas within massive, spherically-symmetric dark-matter halos and compare them with observations of cool-core clusters.  While the outer regions of clusters are certainly not in a steady state because of ongoing accretion and merger events, the inner regions are likely to be closer to a steady state because the relevant timescales for sound waves, gravitational interactions, and radiative cooling are all $\\ll 10$~Gyr at $r \\ll$~100~kpc.  Our exploration proceeds as follows.  Section~\\ref{sec-heq} sets the stage by presenting some useful expressions for the structure of clusters in hydrostatic equilibrium and demonstrating that the temperature profiles of cool-core clusters are determined primarily by the shape of the underlying gravitational potential, as long as the entropy of the central gas is sufficiently low.  Section~\\ref{sec-purecool} then looks at the entropy profiles expected of cluster cores in which cooling is uncompensated by feedback.  These pure-cooling profiles turn out to be nearly power-law in form, $K(r) \\propto r^\\alpha$ with $\\alpha \\sim 1.2$ at large radii.  However, the implied cooling rates of clusters in these configurations are well known to be far larger than the star-formation rates seen in BCGs.  Some form of heating must be compensating for radiative cooling, and we would like to know how the entropy profiles of clusters respond to that heating mechanism.  Section~\\ref{sec-condeq} brings conduction into the mix.  Conduction in the context of cool-core clusters has a long and checkered history.  It has been proposed several times as a mechanism with the potential to balance radiative cooling in cluster cores \\citep{tr83,bm86,nm01}.  However, a cluster in which electron thermal conduction balances radiative cooling is in an unstable configuration \\citep{bd88,Soker03,gor08}.  Furthermore, some cluster cores have cooling rates too large to be balanced by electron thermal conduction, even if it operates at its maximum rate, unfettered by magnetic fields \\citep{Voigt+02,zn03}.  The inability of conduction to suppress cooling flows has led some authors to argue that it is unimportant in cluster cores \\citep[e.g.,][]{Soker10_Bfields}, and there are certainly sharp cold fronts in cluster cores across which conduction must be highly suppressed \\citep{mv07}, presumably by magnetic fields running parallel to the front.  But magnetic fields cannot suppress conduction simultaneously in all directions, and recent work has shown that MHD instabilities driven by anisotropic conduction may actively sculpt the magnetic field geometry in cluster cores \\citep{psl08,pq08,McCourt+10}.  This paper does not investigate the fascinating astrophysics that can result from anisotropic conduction.  Instead it adopts a phenomenological approach in which conductive heat flow runs parallel to temperature gradients, modulo a scalar suppression factor $f_c$ to account for the presence of magnetic fields.  If turbulent velocities in the intracluster medium are sufficiently large, they will overpower buoyancy-driven MHD instabilities and randomize the magnetic field, leading to a mean local suppression factor $\\langle f_c \\rangle \\approx 1/3$ relative to a randomly chosen direction \\citep{pqs10,ro10}.  We will therefore adopt $f_c = 1/3$ as a fiducial value for the scalar suppression factor.  Even though thermal conduction alone cannot stably balance radiative cooling, it remains interesting because it might be governing the triggering of AGN feedback in cluster cores.  \\citet{Donahue+05} raised this possibility in an analysis of clusters with central cooling times much less than a Hubble time but with no obvious evidence for central star formation or AGN feedback.  Those clusters had core entropy levels $\\sim 30-50 \\, {\\rm keV \\, cm^2}$, corresponding to cooling times $\\sim 1$~Gyr, which could in principle be offset by thermal conduction.  \\citet{Cavagnolo+08,Cavagnolo+09} confirmed this relationship between core entropy and feedback in a much larger sample:  Clusters with core entropy $\\lesssim 30 \\, {\\rm keV \\, cm^2}$ often show evidence for AGN feedback while those with higher core entropy levels generally do not.  There are some clusters and groups with powerful central radio sources that appear to have core entropy levels exceeding $\\sim 30 \\, {\\rm keV \\, cm^2}$, but closer inspection reveals that these apparent exceptions have coronae of denser, lower-entropy gas in a kiloparsec-scale region surrounding the central AGN \\citep{Sun09}.  Several studies have shown that thermal conduction, perhaps in concert with AGN feedback, can plausibly explain the observed bimodality in core behavior \\citep{Voit+08,go09,sop09,pqs10,ro10}.  Objects with core entropy $\\lesssim 30 \\, {\\rm keV \\, cm^2}$ have central cooling rates that cannot be balanced by conduction, and AGN feedback is therefore necessary to stabilize cooling.  In objects with greater core entropy, conduction is more efficient than cooling, as long as $f_c \\gtrsim 0.2$, and a modest amount of intracluster turbulence may be necessary to randomize the magnetic fields and keep $f_c$ sufficiently large \\citep{pqs10}.   Feedback from an AGN is not essential to counteract cooling in these objects because of the lower gas density in their cores, and both conductive and dynamical heating are capable of boosting the core entropy until the central cooling time exceeds several Gyr \\citep{co08,pqs10}.  Section~\\ref{sec-condeq} addresses these issues by solving the steady-state equations for intracluster-medium configurations in which electron thermal conduction balances radiative cooling.  While conductively balanced solutions for clusters can differ greatly in density and temperature, they are restricted to a narrow locus in the $K$--$r$ plane and can be approximated with a simple power-law expression. This locus is insensitive to cluster temperature for objects warmer than $\\sim 3$~keV but depends somewhat on $f_c$.  Comparing the locus of conductive balance with the entropy profiles of actual clusters shows that the entropy slopes of cool-core clusters tend to break when they cross the conductively-balanced locus.  At large radii, the $K(r)$ profiles of all clusters lie above the conductively balanced locus and have slopes consistent with gravitational structure formation.  At small radii, the entropy profiles of many cool-core clusters lie at or slightly below the critical locus for conductive balance, with slopes that track the conductively balanced configuration.  This behavior suggests that AGN feedback switches on when thermal conduction can no longer compensate for radiative cooling.  Section~\\ref{sec-thermfeedback} adds the element of thermal feedback.  When conduction is present, steady-state solutions with inward accretion are not allowed unless cooling can outpace conductive heating.  This section shows how steady-state solutions including AGN heating modify the conductively-balanced configuration, allowing {\\em lower} entropy levels to persist in regions where AGN heating exceeds conductive heating.  It also compares steady-state models with cooling, conduction, and thermal feedback with observations of real clusters. These comparisons show that conduction may be a more important heat source than AGN heating in the cores of many cool-core clusters and that AGN heating is usually necessary to counterbalance cooling within only the central $\\sim 30$~kpc.  Section~\\ref{sec-summary} concludes by summarizing these results. ",
        "conclusions": "\\label{sec-summary}  Inspired by the fact that the radial profiles of intracluster gas density, temperature and entropy in cool-core clusters bear a strong resemblance to one another, this paper has tried to identify the physical processes that govern the form of those profiles.  Because the similarity among these profiles extends to within $\\sim 10$~kpc of the cluster's center, where cooling, heating, and dynamical processes can operate on timescales $\\lesssim 100$~Myr, we have assumed that cool-core clusters have settled into a quasi-steady state that is close to hydrostatic equilibrium.   Proceeding from that assumption, the paper derived the general characteristics of various quasi-steady configurations for the intracluster medium.  First, it showed in \\S~2 that the temperature profiles of cool-core clusters simply reflect the shape of the underlying gravitational potential, as long as the intracluster entropy profile has a moderately positive radial gradient.  When that is the case, the gas temperature at each radius is determined primarily by the value of $T_\\phi \\propto M(<r)/r$ at slightly larger radii and to a lesser degree by the slope of the entropy profiles.  The temperature does {\\em not} depend on the absolute value of gas entropy---only the slope of the entropy profile matters.  For an NFW-like potential with a concentration $c_\\Delta \\sim 4$, the gas temperature peaks in the $0.1 r_\\Delta$---$0.2 r_\\Delta$ range, and within a given potential well, configurations with shallower entropy slopes peak at greater temperatures and smaller radii.  Adding a small isothermal potential at the center to represent the mass of a central cluster galaxy helps to explain why clusters contain so little gas at temperatures less than 30\\% of the virial temperature.  Next, \\S~3 derived a set of pure-cooling solutions in which inward advection of heat balances radiative cooling.  Steady cooling of gas in a cluster potential turns out to produce entropy profiles similar in slope to the $K(r) \\propto r^{1.2}$ profile generated by gravitational structure formation.  This coincidence accounts for why there is no break in slope in the entropy profiles of cool-core clusters at radii where the cooling time is similar to the age of the universe.  Furthermore, the mass-inflow rates of these steady pure-cooling solutions are $\\propto K^{-3}$ and greatly exceed the observed star-formation rates of central galaxies in cool-core clusters.  This implies that entropy profiles in cool-core clusters should never drop far below the critical profile at which inflow begins to fuel feedback.  Once inflow begins, the feedback response should rapidly rise as the entropy level at a given radius drops.  Section~4 derived the characteristics of solutions that are in conductive balance, showing that despite widely differing density and temperature profiles, steady configurations with the same value of $f_c$ trace the same locus in the $K(r)$--$r$ plane.  Simple expressions for this locus are given for hot systems ($> 3$~keV) and for cooler systems in which the cooling rate is metallicity dependent.  Comparing the observed entropy profiles of galaxy clusters to these loci shows that clusters without star-forming central galaxies or other evidence of multiphase gas all lie above the locus for conductive balance, suggesting that conduction is preventing radiative cooling from producing a multiphase intracluster medium that can accrete into the central galaxy \\citep[see also][]{Voit+08}.  On the other hand, the entropy profiles of classic cool-core clusters tend to drop below this critical locus at small radii, but they do not drop very far below it.  Instead, their entropy slopes bend in the vicinity of the critical locus and tend to track it at slightly lower entropy levels, suggesting that AGN feedback is triggered when an entropy profile drops below the critical locus and prevents the entropy from dropping much further. This would explain why the density, temperature, and entropy profiles of cool-core clusters are so similar to one another.  Section~5 considered cluster configurations in which cooling, conduction, and thermal feedback all have roles to play in maintaining a steady state.  Instead of assuming a particular model for thermal feedback, the paper attempted to infer the spatial distribution of feedback heating from observations of cool-core clusters.  These were fit to a five-parameter model in which the gravitational potential has an NFW form and the intracluster entropy profile is a power law plus a constant.  This model provides a smooth temperature profile that can be differentiated to obtain the heat deposited by thermal conduction as a function of radius.  If heat advection is assumed to be negligible, the steady-state heating rate is then equal to the difference between radiative cooling and conductive heating.  Some conductive cool-core clusters do not require thermal feedback to offset cooling, but most of them do.  However, the inferred thermal feedback rates usually exceed conductive heating only within $\\lesssim 30$~kpc, where X-ray cavities are most often found.  Conduction with $f_c = 1/3$ can account for most of the heating required beyond that radius, and when integrated over the entire core out to 100~kpc, conduction is a more important heat source than thermal feedback in most of the clusters considered here.  If conduction is indeed as important as these results imply, then it is essential to include it in numerical simulations that seek to accurately model the core structures of galaxy clusters, including the stellar content of the brightest cluster galaxy.   \\vspace*{1.0em}  The author thanks Peng Oh, Ken Cavagnolo, and Marcus Br\\\"uggen for helpful comments on the manuscript.  This work was partially supported at MSU by NSF grant AST-0908819 and {\\em Chandra} Theory program grant TM9-0008X.  It was completed at the Kavli Institute for Theoretical Physics in Santa Barbara, supported in part by the National Science Foundation under Grant No. NSF PHY05-51164."
    },
    "1107/1107.1606_arXiv.txt": {
        "abstract": "At very low frequencies, the new pan-European radio telescope LOFAR is opening the last unexplored window of the electromagnetic spectrum for astrophysical studies. The revolutionary APERTIF phased arrays that are about to be installed on the Westerbork radio telescope (WSRT) will dramatically increase the survey speed for the WSRT. Combined surveys with these two facilities will deeply chart the northern sky over almost two decades in radio frequency from $\\sim 15$ up to 1400 MHz.  Here we briefly describe some of the capabilities of these new facilities and what radio surveys are planned to study fun-damental issues related the formation and evolution of galaxies and clusters of galaxies. In the second part we briefly review some recent observational results directly showing that diffuse radio emission in clusters traces shocks due to cluster mergers. As these diffuse radio sources are relatively bright at low frequencies, LOFAR should be able to detect thousands of such sources up to the epoch of cluster formation. This will allow addressing many question about the origin and evolution of shocks and magnetic fields in clusters. At the end we briefly review some of the first and very preliminary LOFAR results on clusters. ",
        "introduction": "At low frequencies, the new pan-European radio telescope LOFAR is opening the last unexplored window of the electromagnetic spectrum for astrophysical studies. The revolutionary APERTIF phased arrays that are about to be installed on the Westerbork radio telescope (WSRT) will dramatically increase the survey speed for the WSRT. The resulting vast area of new observational parameter space will be fully exploited for many studies directly related to the formation of massive black holes, galaxies, and clusters. Particularly important are three research areas that are driving the design of several surveys that are planned to be carried out with these new facilities. These areas are: (i) forming massive galaxies at the epoch of reionisation, (ii) magnetic fields and shocked hot gas associated with the first bound clusters of galaxies, and (iii) star formation processes in distant galaxies. Furthermore, a most exciting aspect of LOFAR is that its enormous instantaneous field of view coupled with its unprecedented sensitivity at low frequencies equips LOFAR for the discovery of new classes of rare extreme-spectrum sources.  In this contribution, we will first briefly describe LOFAR and APERTIF. For a more extended description of LOFAR we refer to the contribution of George Heald that extensively describes LOFAR and the way the data will be handled to form deep images at low frequencies. Second, we will discuss how  the prime science drivers led to the definition of the planned continuum surveys with LOFAR and APERTIF. Third, we briefly review some of the work we have been carrying out to understand diffuse radio emission associated with merging clusters.  We will mainly concentrate on some of the statistical results obtained for a  partly new sample of relics showing correlations between their sizes, spectral indices and distances from the clusters centres. In the contribution of Reinout van Weeren, he will high-light recent results for newly discovered individual relics, including the recently discovered spectacular double relics in the cluster CIZA J2252.8+5301 (van Weeren et al. 2010). \\nocite{wee10a} Finally, a few preliminary results from LOFAR observations mainly related to clusters are briefly presented. These results show the enormous potential that LOFAR has for studying shocks and magnetic fields in clusters. ",
        "conclusions": ""
    },
    "1107/1107.2542.txt": {
        "abstract": "We work out an effective theory of accelerated expansion for the background to describe general phenomena of inflation and acceleration (dark energy) in the Universe. Our aim is to determine from theoretical grounds,  in a physically-motivated and  model independent way, which and how many (free) parameters are needed to broadly capture the physics of a theory describing the background of cosmic acceleration. Our goal is to make as much as possible transparent the physical interpretation of the parameters describing the expansion. We show that, at leading order, there are five independent parameters, of which one can be constrained via general relativity tests. The other four parameters need to be determined by observing and measuring the cosmic expansion rate only, $H(z)$. Therefore we suggest that future cosmology surveys focus on obtaining an accurate as possible measurement of $H(z)$ to constrain the nature of accelerated expansion. ",
        "introduction": "It has now been experimentally established \\cite{SN,SNII,LSS,Stern,Jimenez} that the Universe is currently accelerating and present knowledge indicates that eventually will enter a deSitter phase (see the reviews in Ref.~\\cite{Peebles:2002gy,Copeland:2006wr,Sahni:2006pa,Frieman:2008sn} and references therein). There is also experimental evidence that large--scale structures in the  early Universe were seeded by non-casual perturbations \\cite{wmap01}; the best theoretical model to explain this is inflation, which consist of an accelerating de-Sitter phase in the  early Universe. Therefore, there is now strong experimental evidence that the Universe has experienced periods of accelerated expansion, however we have no satisfactory explanation of what has been driving it and of the physical mechanism behind it.  Progress can arise both from the theoretical and observational front. It can happen that a fundamental theory is found that tells us what the nature of inflation and/or dark energy are. On the other hand, it could be that we need to exploit astronomical observations in order to zoom in into the detailed properties of the expansion before we can discriminate among competing theories and shed light on the mechanism driving the expansion.  If  the latter scenario is realized the question that arises is: how can we determine in the most model-independent way, the nature of acceleration? Can this be done in such a way that observational constraints can have an easy and transparent interpretation in terms of the properties of a possible underling physical mechanism? For example, in the context of dark energy, the most widely used approach is to assume that the expansion is driven by a scalar field with constant equation of state $p = w \\rho$ and use observations to find $w$ ($p, \\rho$ are the pressure and energy density of the scalar field). Of course, in this framework, a cosmological constant corresponds to $w=-1$ while scalar fields with dynamics deviate from this value. It is clear that one cannot assume $w$ to be constant with time, and that a more general scenario would require a reconstruction of $w(t)$. It is important to keep in mind that this description is a drastic simplification as it  does not cover all possible scenarios (see e.g.,\\cite{Simon}), and  there could be  models where $p, \\rho$ are not even defined,  but from observations of any expansion history  one could always reconstruct  an effective $w(t)$. In this case the $w(t)$ would have no physical meaning in terms of properties of a fluid or a scalar field and would be of difficult or ambiguous  theoretical interpretation. Even with this drastic simplification, it is very challenging observationally to determine $w(t)$ \\cite{Simon}. To circumvent this problem, the community has proposed several parameterizations: some more closely related to the underlying physics (e.g.,\\cite{Simon}) and others purely phenomenological (\\cite{Polarski,Linder} and refs therein). All these parameterizations make compromises in order to adjust the number of parameters to the observables in the sky. Even if we do not consider this limitation, it is unclear how closely related they are to tell us something about the underlying physics driving the expansion. Moreover degeneracies can be found such that a constant equation of state can be mimicked by a variable one in these parameterizations (e.g.,\\cite{JimenezLoeb2002}).  Given the above limitations, it is worth asking the question of how to build a general theory of accelerated expansion that captures as much as possible the physics and thus avoids arbitrary parameterizations. This is what we set-up to do in this paper: we build an effective theory for the background of cosmological accelerated expansion. While the philosophy of our approach can be recognized in the works of e.g., Ref~\\cite{Weinberg:2008hq,Creminelli,Park:2010cw} we note that our focus is to deal with the background and not with the fluctuations, as was done by the previously referred papers which in turn were inspired by the theory of pions from the 70's. A similar approach has been developed in \\cite{Park:2010cw} where the authors have studied thoroughly the phenomenological consequences of their approach. In our approach we propose a physically--motivated way to truncate the effective theory of expansion when there is no physical information about the driver of acceleration, e.g. we do not want to constraint our model to be a quintessence one. In inflation one observes directly the fluctuations but in dark energy we only see the expansion, thus we expect our methodology to be more relevant in the context of dark energy. ",
        "conclusions": "We have developed an effective theory of the background of accelerated expansion with the aim of  making a direct connection  between  cosmological observables and  parameters  describing  the physics behind the expansion. We have adopted what we propose as the ``natural\" power counting to expand (and truncate)  the expression for the Lagrangian describing the accelerating Universe. Our choice is motivated by being natural for scalar fields.  In doing so we have discovered that only five parameters are needed to fully described the (effective) theory at leading order. These parameters are readily connected to the physics behind expansion and thus are the most natural set an observer should  attempt to determine. In particular one parameter  describes deviations from standard --general relativity-- gravity and the other  four completely determine the expansion history. Although our treatment is general to both inflation   and dark energy, we have concentrated on the case for dark energy which yields less standard results.  Our conclusion is  that one should do observations that test deviations from general relativity (via e.g., constraints on the growth of cosmological structure or higher-order correlations of the matter  density field), and measure the Hubble parameter $H(z)$ with  the best possible  accuracy in at least four redshift bins. Once $H(z)$ has been measured, Eq.~ (\\ref{eq:master}) can be used to obtain the physical parameters of the Lagrangian that describes the accelerated expansion.  Note that, even neglecting modifications to General Relativity, if from observations one wanted to  reconstruct in a non-parametric way the potential of the accelerating field, both the Hubble parameter  $H(z)$ and is derivative $dH(z)/dt$ would be needed as shown in \\cite{Simon}. Here, the effective theory expansion already provides a parameterization of the potential as a function of the field, and thus only a measurement of $H(z)$ and  the matter density are needed.  A priori we have no restrictions on the actual values of the unknowns in the theory, beside imposing the validity of the effective theory expansion. If, as an outcome of  confronting the theory with  experimental data in Eq.~ (\\ref{eq:master}) some of these parameters turns to be numerically suppressed, we would have learned something about  the pattern of the explicit symmetry breaking for the scalar field. If, on the contrary,  it turns out their numerical values do not match the a priori power--counting estimates,  this would be an indication of a break down of the proposed approach and that an alternative power--counting must be implemented. Although we have tackled only the leading contribution, perturbations on any of both fields can be computed in a straightforward manner  along  the lines of Ref.~\\cite{Weinberg:2008hq}. For instance, the gravitational corrections turns to be identical to those found by Ref.~\\cite{Weinberg:2008hq}. In our case the speed of sound equals unity, $c_s=1$. Corrections to this value come from higher order operators as in Eq.~(\\ref{supp}).  Future cosmology surveys promise to provide measurements of the expansion rate --either directly measuring H(z), or closely related quantities such as the angular  diameter distance  or the luminosity distance as a function of redshift-- with percent precision over a wide range of redshift ($z \\sim 3$)  and over tens of redshift bins. This will offer a unique opportunity to gain insight   into the mechanism of cosmic acceleration. In recent work \\cite{Jimenez:2012jg} we have used the newly determined $H(z)$ data by Ref.~\\cite{Moresco:2012jh} to put constraints on the effective potential of dark energy and found it to be consistent with a flat potential, with an allowed deviation of 6\\% from flat (Fig.~3 in \\cite{Jimenez:2012jg}). Further, we showed how the field has moved in most of the models constraints, thus illustrating the power of the effective field theory approach to elucidate the nature of accelerated expansion: if the fact that the field has moved is confirmed by future and more accurate measurements of $H(z)$, this would indicate that dark energy was caused by a  pseudo-goldstone boson and therefore it was the result of a symmetry breaking. In this paper we have concentrated on the background expansion and neglected perturbations. While this is satisfactory for dark energy, as the observational perspectives of detecting dark energy fluctutations are very gloomy, it is not for inflation, where perturbations are relevant."
    },
    "1107/1107.3603_arXiv.txt": {
        "abstract": "To investigate the physics of mass accretion onto weakly-magnetized neutron stars, {95} archival {\\it RXTE} data {sets} of an atoll source 4U 1608--522, {acquired over 1996-2004} in so-called upper-banana state, were analyzed. {The object meantime exhibited 3--30 keV luminosity in the range of $\\lesssim 10^{35} - 4 \\times 10^{37}$ erg s$^{-1}$, assuming a distance of 3.6 kpc.} {The 3--30 keV PCA  spectra, produced one from each dataset,} were represented successfully with a combination of a soft and a hard component, {of which the presence  was revealed in a model-independent manner by studying spectral variations among the observations.} The {soft component} is expressed by so-called multi-color disk model with a temperature of $\\sim 1.8$ keV, and is attributed to the emission from an optically-thick standard accretion disk. The {hard component} is a blackbody emission {with a temperature of} $\\sim 2.7$ keV, thought to be emitted from the neutron-star surface. As the total luminosity increases, a continuous decrease was observed in the ratio of the blackbody luminosity to that of the disk component. This property suggests that the matter flowing through the accretion disk gradually becomes difficult to reach the neutron-star surface, presumably forming outflows driven by the increased radiation pressure. {On time scales of hours to days, the overall source variability was found to be controlled by two independent variables; the mass accretion rate, and  the innermost disk radius which changes both physically and artificially.} ",
        "introduction": "It has long been known that  low-mass X-ray binaries (LMXBs), namely close binaries involving neutron stars (NSs) without strong magnetic fields, {emit thermal type spectra when they are luminous ($\\gtrsim 10^{36}$ erg s$^{-1}$)}. In the 1970's, {their spectra} were often modeled empirically in terms of simple thermal Bremsstrahlung, without much physical basis. As observations made progress, it gradually became clear that this simple empirical modeling is in fact inadequate, and at least two components are needed to describe \\green{wide-band (typically 2--30 keV) LMXB spectra. Today, a  canonical  model (sometimes called ``Eastern model'') interprets the X-ray spectra of  a luminous LMXB as originating from two major emission regions; the surface of the NS \\green{(or often called ``boundary layer'')} and an accretion disk around it \\citep{mitsuda_z}.}  The Eastern model was \\green{constructed} based on {\\it Tenma} observations of luminous LMXBs \\citep{mitsuda_z,mitsuda_1608,makishima_gx3+1}, particularly considering intensity-correlated spectral changes. {When an LMXB becomes brighter on a time scale of $\\sim$ 1000 s, its spectrum  hardens in energies from $\\sim 2$ to $\\sim 10$ keV,} but the spectral shape in the hardest range (above 10 keV) stays constant. A difference between a pair of spectra with different intensities generally {takes a form of \\red{a blackbody (BB) \\blue{spectrum}}, of which} the intensity changes but the shape (i.e., the temperature) is approximately constant. {This BB component has a temperature of $\\sim$ 2 keV, which} is close to the Eddington temperature, 2.0 keV, for a NS of $\\sim$ 1.4 $M_{\\odot}$ mass (with $M_{\\odot}$ being the solar mass) \\blue{and 10 km radius.} {In addition, the area of its emission region is inferred to be a fraction of the surface area of \\blue{such a NS}.} Therefore, the BB component has been ascribed successfully to the NS surface emission. Indeed, every spectrum observed from these sources was reproduced by a sum of this BB component, and an additional soft component. The soft component was approximated first by a softer BB of $\\sim 1$ keV temperature, but later, much better by a particular superposition of BB spectra, called ``disk blackbody\" or ``multi-color disk\" (MCD) emission \\citep{mitsuda_z,makishima_gx3+1}, as predicted by the optically-thick standard accretion disk model \\citep{shakura_disk}.  While the Eastern model thus provides a promising ground to understand  the physics of mass accretion {in luminous LMXBs}, their complex spectral and intensity variations have often been described in a purely empirical manner using ``color-color'' diagrams and other similar methods. As a result, such empirical classifications as ``Z sources'' and ``atoll sources'' have been created, together with various ``branches'' on these empirical diagrams \\citep[e.g., ][]{hasinger_z}. These primitive descriptions are waiting to be replaced by more physically meaningful ones, using \\blue{the Eastern model}.  After the launch in 1995, the {\\em Rossi X-Ray Timing  Explorer} % \\citep[{\\it RXTE};][]{bradt_rxte} has been observing many LMXBs for a huge number of times. With the largest effective area and the highest timing resolution ever achieved, {\\it RXTE} is {indeed} very suitable to the study of {spectral variations} of LMXBs. {On time scales of milliseconds to seconds, \\citet{gilfanov_vari} and \\citet{revnivtsev_vari} carried out such studies using the {\\it RXTE} data, and revealed that the variation is carried by a rather hard  BB-like emission component with a temperature of $\\sim$ 2--3 keV.} These results reinforce the validity of the Eastern model, and encourage us to attempt its extensive application to the {\\it RXTE} data of LMXBs. {We thus} hope to understand the spectral variations of LMXBs as a whole using a physical model (based on the Eastern description).   \\if Contrary to the Eastern model, the Western model assumes that Thomson opacity dominates free-free opacity in the accretion disk and the emission is significantly modified by electron scattering, although the emission from the NS surface are hardly modified and observed as simple BB emission. \\citet{white_west} for the first time applied this model to actual X-ray data, and argued that it can reproduce the observed LMXB spectra best among many models they attempted, including the Eastern model without the consideration of the Comptonization hard tail. As a variant to this model, there have been repeated suggestions that the accretion disk is surrounded by a cloud of hot electrons, or ``accretion disk corona'', which may add to the Comptonization. \\citet{church_west} analyzed dip sources, and also argued that the Western model can reproduce their spectra in a unified way at various dipping levels, and hence with different shapes. \\fi   As the first of a series of {our planned} publications, the present paper deals with so-called upper-banana state {of 4U 1608--522, which is one of  the most frequently observed  atoll sources}  by {\\it RXTE}. {Section 2 describes  414  {\\it RXTE} data sets of  this LMXB, of which 95 are selected for use in the present paper. In \\S~3.1, we select four representative spectra out of the 95 data sets, and  analyze  their intensity-correlated spectral changes to reconstruct the Eastern model. The derived model parameter values are examined in \\S~3.2, followed by \\S~3.3 where the model  is applied  to all the 95 energy spectra and  relations among the obtained physical parameters are studied. As a reconfirmation of our approach, we calculate in \\S~3.4 the effective degrees of freedom involved in the variability, of which the results are physically interpreted in \\S~3.5. As discussed in \\S~4, the results} give a full support to the Eastern model, and lead to a finding of mass outflows as the source gets luminous. ",
        "conclusions": "\\subsection{{The Eastern Model}}  Through the analysis of the {\\it RXTE} spectra of the atoll source 4U 1608--522 in its upper-banana state, we have confirmed that both time-averaged and difference energy spectra can be reproduced successfully by the Eastern (MCD+BB+Gau) model. Considering the hardening factor, the effective temperature and radius of the MCD component were obtained as $kT^{\\rm eff}_{\\rm in} \\sim$ 1.1 keV and $r^{\\rm eff}_{\\rm in} \\sim$ 12 km, respectively. The radius is larger than the representative NS radius, 10 km, and is consistent with the last stable orbit, $3 R_{\\rm s}$ = 12.4 km, allowed by general relativity. This result agrees with the picture of a standard accretion disk which is formed around a NS. The BB parameters were found as $kT^{\\rm eff}_{\\rm BB} \\sim$ 1.8 keV and $r^{\\rm eff}_{\\rm BB} \\sim$ 2.7 km, assuming isotropic emission from a spherical source. The temperature is close to the local Eddington temperature (2.0 keV) at the NS surface, and the radius is smaller than 10 km. Then, the BB component can be regarded as being emitted from an equatorial zone of the NS, as previously suggested by \\citet{mitsuda_z}.   \\subsection{Saturation of the Blackbody Luminosity}  According to the picture of standard accretion disks, a half of the released gravitational energy is radiated from the accretion disk ($L_{\\rm disk}$), and the other half is stored in the Keplerian kinetic energy and then emitted ($L_{\\rm BB}$) when the matter settles onto the NS surface. Then, we expect $L_{\\rm BB}$ to be proportion to $L_{\\rm disk}$, and hence the $L_{\\rm BB}/L_{\\rm disk}$ ratio to be constant. Indeed, $L_{\\rm disk}$ was found to increase almost linearly with the total luminosity $L_{\\rm tot}$ (figure 6a). However, the same figure reveals that $L_{\\rm BB}$ increases less steeply, making the $L_{\\rm BB}/L_{\\rm disk}$ ratio decrease from 0.6 to 0.4 as $L_{\\rm tot}$ increases from $1 \\times 10^{37}$ to $4 \\times 10^{37}$ erg s$^{-1}$.   We may think of two possibilities to explain the observed relative decrease of $L_{\\rm BB}$. One is that the accreting matter reaches the NS surface and emits $L_{\\rm BB}$ which is equivalent to the Keplerian energy (i.e. $\\propto L_{\\rm disk}$), but we cannot observe all the emission because its wavelength shifts outside the PCA energy band (3--30 keV), or the geometrical angle of the emission moves from our line of sight. However, we have not observed any hint of extra emission in the softest or hardest energy ends of the PCA spectra. In addition, other sources, which are thought to have different inclination angles, are also reported to show the same behavior of the decreasing $L_{\\rm BB}/L_{\\rm disk}$ ratio \\citep{revnivtsev_vari}. Therefore, this possibility is unlikely.  The other possibility is that the accreting matter flows all the way through the disk down to its inner radius that is close to the NS surface, but its progressively larger fraction fails to accrete onto the NS surface. If the accretion-failed matter stayed around the NS away from the surface (e.g., an expansion of the boundary layer), it would accumulate to become optically thick, eventually producing detectable emission. This would lead to a decrease in $kT_{\\rm BB}$, because the emission radius should increase. As shown in Figure~\\ref{fig:spec_1608_all1} and equation~(\\ref{eq:1608_Ltot_t_r}), however, this is not the case; the $kT_{\\rm BB}$ value stays almost constant (or rather increases) as the total luminosity increases. Then, a fraction of the matter must be outflowing and escaping from the system without releasing its kinetic energy as emission. \\blue{Although the observed $L_{\\rm tot} \\sim 4 \\times 10^{37}$ erg s$^{-1}$ is only $\\sim 20\\%$ of the  Eddington luminosity for a 1.4 $M_{\\odot}$ NS, the observed value of $kT^{\\rm eff}_{\\rm BB} \\sim 1.8$ keV is already close to the spherical Eddington temperature at 10 km (2.0 keV). Considering further the non-spherical geometry of the disk, and the disk radius which is possibly larger than 10 km, the putative outflow is likely to be driven by increased radiation pressure in the innermost disk region. } The presence of such outflows is indeed suggested by detections of broad absorption features from LMXBs (Schulz \\& Brandt 2002; Ueda et al.\\ 2004), black hole binaries (Kotani et al.\\ 2000; Yamaoka et al.\\ 2001; Kubota et al.\\ 2007), and  active galactic nuclei (Pounds et al.\\ 2003ab).  Assuming that the decrease in  $L_{\\rm BB}/L_{\\rm disk}$ from 0.6 to 0.4 is due to outflows, about $(0.6-0.4)/0.6 \\sim$ 30\\% of the accreting matter is estimated to escape from the system at a typical luminosity of $\\sim 4 \\times 10^{37}$ erg s$^{-1}$. \\green{Toward lower luminosities, the ratio appears to converge to $\\sim 0.6$, rather than 1.0 which would be theoretically expected. This may be attributed to effects due, e.g., to the system inclination, general relativity, and NS rotation (e.g., \\citet{sunyaev_ratio}). Detailed  theoretical discussion on this point is  beyond the scope of the present paper.}    As indicated by equation~(\\ref{eq:1608_Mdot_disk2}), the innermost disk radius $r_{\\rm in}$ is observed to increase gradually \\green{when the mass accretion rate increases to$\\gtrsim 300$ in Figure~\\ref{fig:spec_1608_all3}, or equivalently, when the total luminosity becomes $ \\gtrsim 1.5 \\times 10^{37}$ erg s$^{-1}$ (corresponding to \\blue{$\\gtrsim 7\\%$} of the Eddington luminosity).} This is unlikely to be an apparent effect caused, e.g., by changes in the hardening factor, since this quantity would increase as the luminosity increases \\citep{gierlinski_hard,davis_hard}, and would hence make $r_{\\rm in}$ smaller. Therefore, the increase in $r_{\\rm in}$ toward higher accretion rates is considered to represent a real retreat of the innermost disk radius. This interpretation \\red{may} agree with the results by \\citet{popham_boundary}, who numerically calculated the physical condition of the boundary layer and showed that the disk inner edge retreats back by the increased radiation pressure when the luminosity approaches the Eddington limit.   \\if In Figure~\\ref{fig:spec_1608_all1} and \\ref{fig:spec_1608_all3}, the increase of $L_{\\rm BB}$ is mainly caused by that of $r_{\\rm BB}$, while $kT_{\\rm BB}$ {remains almost constant.} This may be also because of the high-radiation pressure of $kT_{\\rm BB}$ close to the Eddington temperature, since the matter tends to escape such high-pressure regions and accrete onto other places, expanding the emitting equatorial region on the NS surface rather than increasing the temperature. \\fi  \\if Along the increase of $\\dot{M}$, we consider that the physical state of the accretion flow will evolve in the following manner. At the beginning of the upper-banana state ($\\dot{M} \\lesssim 300$), there should be little outflows, and all the matter accretes from the standard accretion disk to a relatively narrow equatorial region on the NS surface. In this stage, the BB emission comes from the narrow belt region of the NS surface. When $\\dot{M}$ increases, the accretion disk behaves in the ``standard'' manner; the innermost temperature $kT_{\\rm in}$ is proportion to $\\sim \\dot{M}^{0.25}$, and the innermost radius $r_{\\rm in}$ stays almost constant. Since $kT_{\\rm BB}$ stays uniform and almost consistent with the Eddington temperature, we then expect the belt region emitting the BB component becomes vertically wider, and hence the observed $r_{\\rm BB}$ increases. \\fi  \\if As $\\dot{M}$ increases $\\gtrsim 300$, the disk may stay in the standard state, but the transition region will become vertically thicker due to the increased radiation pressure inside the disk. Due to the temperature uniformity, the outward radiation pressure is expected to be latitudinally rather constant, while the inward gravitational pressure must decrease toward higher (positive and negative) latitudes due to the reduced matter density. As a result, we expect the outflow to take place first at the ``northern'' and ``southern'' edge of the belt, and proceed toward the equatorial region. This scenario presented above explains, at least qualitatively, the relative decrease in the $L_{\\rm BB}/L_{\\rm disk}$ (and hence $L_{\\rm BB}/\\dot{M}$) ratio. \\fi   \\subsection{{Secondary Spectral Variations}}  The result of the fractal dimension analysis shows that the source variations are controlled by two independent variables. One of them is obviously the mass accretion rate (or nearly equivalently the total luminosity), which produces the major trend in Figure~\\ref{fig:spec_1608_all3} as discussed so far. The other variable is thought to be random changes in the inner disk radius, as represented by the behavior of $r'_{\\rm in}$  in the de-trended analysis performed in \\S~\\ref{subsec:vari}. As discussed therein, the variations of $r'_{\\rm in}$ may be caused by two different mechanisms, corresponding to the CLB and VLB. These two branches are considered to degenerate in the fractal dimension analysis, and difficult to identify as two independent degrees of freedom.   In the CLB, the change in $r'_{\\rm in}$ seems to be apparent rather than real, and is considered to reflect fluctuations in the hardening factor of the disk As a result, the BB parameters remain unchanged, while the MCD parameters vary as  $kT'_{\\rm in} \\propto r'_{\\rm in}$$^{-0.5}$, keeping the disk luminosity constant ($L'_{\\rm disk} \\propto r'_{\\rm in}$$^{0}$). Such variations in the hardening factor may in turn stem, e.g.,  from fluctuations of the disk height, because this would affect the temperature gradient and/or radiative transfer in the vertical direction of the disk.   In the VLB, the change in $r'_{\\rm in}$ is considered real (\\S~\\ref{subsec:vari}), because we observe a scaling of $kT'_{\\rm in} \\propto r'_{\\rm in}$$^{-0.75}$ which agrees with the behavior of a standard accretion disk under a constant $\\dot{M}$ and a constant hardening factor. In this branch, the disk is considered to randomly retreat back from its ``home position\", which is located in Figure~\\ref{fig:detrend} around $r'_{\\rm in} \\sim 0.85$ where the CLB and VLB meet. When this takes place, the accreting matter located between the retreated and the home-position radii will fall onto the NS surface without radiating the disk emission, and the gravitational energy it acquires will be released solely as the radiation from the NS surface. Then, $L_{\\rm disk}$ will decrease while $L_{\\rm BB}$ will increase. This can naturally explain the anti correlation between the MCD and BB luminosities observed in the VLB (Figure~\\ref{fig:detrend}).    \\subsection{Conclusion} The results of our study can be summarized in the following four points. \\begin{enumerate} \\item The physical picture provided by the Eastern model can consistently explain the behavior of 4U 1608--522 in the upper-banana state. \\item When the mass accretion rate increases, some fraction of the matter flowing through the accretion disk starts outflowing, presumably due to an increased radiation pressure. \\item Even at the same accretion rate, the spectral parameter changes randomly on time scales of hours to days, mainly reflecting changes in the innermost disk radius. \\item The change in the disk inner radius is a mixture of real and apparent effects, the former due to sporadic retreat of the disk while the latter due to fluctuations in the color hardening factor. \\end{enumerate}    Facilities: \\facility{RXTE (PCA)}."
    },
    "1107/1107.2605_arXiv.txt": {
        "abstract": "Previously a theory has been presented which extends the geometrical structure of a real four-dimensional space-time via a field of orthonormal tetrads with an enlarged transformation group.  This new transformation group, called the conservation group, contains the group of diffeomorphisms as a proper subgroup. Field equations were obtained from a variational principle which is invariant under the conservation group.  In this paper, this theory is further extended by development of a suitable Lagrangian for a field with sources. Spherically symmetric solutions for both the free field and the field with sources are given.  A stellar model and an external, free-field model are developed. The theory implies that the external stress-energy tensor has non-compact support and hence may give the geometrical foundation for dark matter.  The resulting models are compared to the internal and external Schwarzschild models.  The theory may explain the Pioneer anomaly and the corona heating problem. ",
        "introduction": "Let ${\\cal X}^4$ be a 4-dimensional space with orthonormal tetrad $h^{i}_{\\,\\,\\mu}$.  Then a metric $g_{\\mu \\nu}$ may be defined on ${\\cal X}^4$ by $g_{\\mu\\nu}=\\eta_{ij}\\, h^i_{\\,\\,\\mu}h^j_{\\,\\,\\nu}$ where $\\eta_{ij} = diag\\bigl\\{-1,1,1,1\\bigr\\}$.  Whereas Einstein extended special relativity to general relativity by extending the group of transformations from the Lorentz group to the group of diffeomorphisms, Einstein later suggested that that a unified field theory may be obtained by extending the diffeomorphisms to a larger group [1].  Einstein was also led by the principle that the speed of light was constant.  Consistent with Einstein's approach, we look for the largest group of transformations for which the wave equation, $\\Psi^\\alpha_{\\;\\; ; \\alpha} = 0$, is covariant.   This is the guiding principle for our theory.  Let $\\tilde{V}^\\alpha$ be a vector density of weight $+1$.  Then a conservation law of the form $\\tilde{V}^\\alpha_{\\;\\; , \\alpha}=0$ is invariant under all transformations satisfying \\begin{equation}x^\\nu_{\\;\\; ,\\overline{\\alpha}}\\bigl(x^{\\overline{\\alpha}}_{\\;\\; ,\\nu ,\\mu} - x^{\\overline{\\alpha}}_{\\;\\; ,\\mu , \\nu} \\bigr) = 0 \\quad .  \\end{equation} This property defines the group of conservative transformations, of which, the group of diffeomorphisms is a proper subgroup [2].  Since the wave equation may be written as $\\tilde{V}^\\alpha_{\\;\\; ,\\alpha} = 0 $ with $\\tilde{V}^\\alpha = \\sqrt{-g}\\Psi^\\alpha$, we see that the conservation group is \"the largest group of coordinate transformations under which the equation for the propagation of light is covariant\" [3].  Although we may view the space ${\\cal X}^4$ as a Riemannian manifold, the space is more general than a manifold [2-6]. Suppose that $x^\\mu$ are coordinates on a Riemannian manifold ${\\cal M}$.  Then for functions $f$ with continuous second derivatives, the commutator of partial derivatives is zero, i.e., $[\\partial_\\mu ,\\partial_\\nu ]f=0$.   If $x^{\\bar{\\alpha}}_{\\;\\; ,\\mu}$ is non-diffeomorphic but satisfies (1) we see that $[\\partial_{\\bar{\\mu}}, \\partial_{\\bar{\\nu}}]f$ is nonzero. Thus the geometry associated with the group of conservative transformations is more general than a manifold.  It is not simply an abandonment of the diffeomorphism condition that $[\\partial_\\mu,\\partial_\\nu] = 0$, however, because we require that the transformation be conservative.  The argument that accelerated observers should be on equal footing led Einstein to general relativity.  We have argued that requiring that quantum observers be on equal footing leads to the conservation group [3].  If we, as observers consider ourselves to be classical (non-quantum) observers, we will have a preference for the manifold view for what we observe. In truth, we are are quantum observers and hence some \"fuzziness\" in our observations as well as the observations of other observers is present.  Suppose $x^\\alpha$ are used as coordinates on a neighborhood of our manifold.   If $x^{\\bar{\\mu}}_{\\; ,\\nu}$ is non-diffeomorphic but satisfies (1), then $x^{\\bar{\\mu}}$ is similar to anholonomic coordinates [7] .  However, the conservative group property (1) ensures that observers, one using $x^\\alpha$ and the other using $x^{\\bar{\\mu}}$ as their coordinates, are on equal footing.  When we take the manifold view, we interpret the change of coordinates from $x^\\alpha$ to $x^{\\bar{\\mu}}$ as only making sense at the differential (i.e. infinitesimal) level.   Thus, we see the neighborhoods upon which the coordinate systems of the second observer make sense to us as the first observer have shrunk from global (special relativity) to local (general relativity) to infinitesimal (conservation group theory).  We stipulate that we may begin the setup of our theory as if it were a manifold, defining $h^i_{\\;\\alpha}$ as a function of $x^\\mu$.  On equal footing there are many coordinate systems related to $x^\\mu$ via  conservative transformations, however we, as observers, have a natural preference for the manifold view and hence prefer $x^\\mu$.  In an effort to model dark matter cosmic acceleration, many theorists have simply modified general relativity in some fashion.  We claim that our modification which is based on an enlargement of  the transformation group is perhaps the only one with a solid guiding principle.  Einstein himself felt it was a mistake to simply add a cosmological constant $\\Lambda$.  Recently, theoretical developments of $f(R)$ gravity [8], quintessence [9] and other modifications of general relativity [10] have a similar {\\it ad hoc} flavor.  The geometrical content of the theory based on the conservation group is determined by $C_\\alpha \\equiv h_i^{\\,\\,\\nu}\\bigl(h^i_{\\,\\, \\alpha ,\\nu}-h^i_{\\,\\, \\nu ,\\alpha} \\bigr) \\; = \\gamma^\\mu_{\\;\\;\\alpha\\mu}$, where the Ricci rotation coefficient is given by $\\gamma^i_{\\;\\;\\mu\\nu}=h^i_{\\;\\mu ;\\nu}$ [2-6]. Pandres calls this the curvature vector. He shows that $C_\\alpha$ is invariant under transformations from $x^\\mu$ to $x^{\\overline{\\mu}}$ if and only if the transformation is conservative and thus satisfies (1).  A suitable scalar Lagrangian for the free field is given by \\begin{equation} {\\cal L}_f= \\frac1{16\\pi}\\int C^\\alpha C_\\alpha \\, h \\; d^4 x  \\end{equation} where $h = \\sqrt{-g}$ is the determinant of the tetrad.  Using $h^i_{\\, \\mu}=h^I_{\\, \\mu}\\Lambda^i_I$, we have extended the field variables [5] to include the tetrad $h^I_{\\;\\mu}$ and 4 internal vectors $\\Lambda^i_I$, with internal space variable $x^I$. With this extension, the covariant derivative has been extended to be invariant under a larger group of transformations on $x^I$ as well as $x^\\mu\\;$ [5].  We assume that the metric on the $x^I$ space is also $\\eta_{IJ}=diag\\bigl(-1,1,1,1\\bigr)$.  The definition of the Ricci rotation coefficient is also extended using the $\\Lambda^i_I$ to \\begin{equation}\\Upsilon^\\alpha_{\\;\\;\\mu\\nu}\\equiv h_I^{\\;\\alpha}h^I_{\\;\\mu ;\\nu}+h_i^{\\;\\alpha}h^I_{\\;\\mu}\\Lambda^i_{I,\\nu} \\end{equation} and the definition of $C_\\alpha$ is also extended to $C_\\alpha \\equiv \\Upsilon^\\mu_{\\;\\;\\alpha\\mu}$. Using these extended Ricci rotation coefficients, one finds that \\begin{equation}C^\\alpha C_\\alpha = R + \\Upsilon^{\\alpha \\beta \\nu} \\Upsilon_{\\alpha \\nu \\beta}-2C^\\alpha_{\\; ;\\alpha} - \\eta^{ij}h_j^{\\;\\nu}h_I^{\\;\\alpha}(\\Lambda^I_{i,\\alpha ,\\nu}-\\Lambda^I_{i,\\nu ,\\alpha})\\quad , \\end{equation} where $R$ is the usual Ricci scalar curvature. Thus, when the physical space is interpreted as a manifold, the Lagrangian density of the free field contains terms corresponding to non-gravitational interactions [4,6].   The motion of a free particle or photon in the inertial coordinate system is given by \\begin{equation} \\frac{d^2 x^i}{ds^2}=0\\, , \\end{equation} where $-ds^2=\\eta_{ij}dx^i dx^j$.  This equation when transformed to internal coordinates, $x^I$ is \\begin{equation} \\frac{d^2x^I}{ds^2}=-\\Lambda^I_i \\Lambda^i_{J,K} \\frac{dx^J}{ds\\;}\\frac{dx^K}{ds\\;} \\; , \\end{equation} where the right hand side of this equation is zero when there are no internal forces.  Since $\\eta_{IJ}$ corresponds to the flat metric, we naturally interpret the right hand side of (6) as a force.   The $\\Lambda^i_I$ are thus internal fields that via $\\Lambda^i_{I,J}$ correspond to electroweak and strong interactions.  In the manifold view, with coordinates $x^\\alpha$  equation (6) becomes \\begin{equation} \\frac{d^2 x^\\alpha}{ds^2}+ \\Gamma^\\alpha_{\\mu\\nu}\\frac{dx^\\mu}{ds\\;}\\frac{dx^\\nu}{ds\\;} = - \\Upsilon^\\alpha_{\\;\\;\\mu\\nu}\\frac{dx^\\mu}{ds\\;}\\frac{dx^\\nu}{ds\\;} \\; . \\end{equation} The right hand is partly generated from the internal forces since from (3) one sees that this equation of motion depends on $\\Lambda^i_{I,\\nu}$.   Setting the variations of ${\\cal L}_f$ with respect to $h^I_{\\;\\;\\mu}$ and $\\Lambda^i_I$ equal to zero along with the assumption that we may always choose $\\Lambda^i_I$ to correspond to a complex Lorentz transformation (since $h^i_{\\;\\mu}=h^I_{\\;\\mu}\\Lambda^i_I$), yields the field equations \\begin{equation}  C_\\mu = 0 \\quad .  \\end{equation}  One feature of the extended theory with field variables $h^I_{\\;\\mu}$ and $\\Lambda^i_I$ is that the internal fields associated with $\\Lambda^i_I$ may be specified after finding a tetrad $h^I_{\\;\\alpha}$ which satisfies the condition $h_I^{\\;\\nu}\\bigl(h^I_{\\; \\mu ,\\nu}- h^I_{\\; \\nu , \\mu}\\bigr)=0$ .  If this tetrad is a function of position, then $h^I_{\\;\\alpha}$ yields a Riemannian manifold (manifold view) with corresponding metric, $g_{\\mu\\nu} = \\eta_{IJ}h^I_{\\;\\mu}h^J_{\\;\\nu}$.  Changes in $\\Lambda^i_I$ have no effect on this manifold [5].  If $\\Lambda^i_I$ is a constant field such that $\\eta_{ij}=\\eta_{IJ}\\Lambda^I_i \\Lambda^J_j = diag \\bigl(-1,1,1,1\\bigr)$, then the field equations are satisfied, however $\\Lambda^i_I$ may be non-constant.  There exist nonconstant (non-diffeomorphic) values of $\\Lambda^i_I$ that satisfy the conservative condition,  $\\Lambda^J_k\\bigl(\\Lambda^k_{I,J}-\\Lambda^k_{J,I}\\bigr)=0$.  With $\\Lambda^i_I$ that satisfy this condition, the field equations,  $C_\\mu = 0$ remain satisfied. {\\it Since this paper is concerned with gravitational implications of the the theory we will assume for the remainder of this paper that we are working with a solution of the field equations for which $\\Lambda^I_i =\\delta^I_i $.}  In this case, an identity for  the Einstein tensor is \\begin{eqnarray}  G_{\\mu\\nu}= & C_{\\mu ; \\nu}- C_\\alpha \\Upsilon^\\alpha_{\\; \\mu\\nu} -g_{\\mu\\nu}C^\\alpha_{\\; ;\\alpha}-\\frac12 g_{\\mu\\nu}C^\\alpha C_\\alpha   \\nonumber \\\\ \\quad & +\\Upsilon^{\\;\\;\\alpha}_{\\mu \\;\\; \\nu ;\\alpha}+\\Upsilon^\\alpha_{\\;\\;\\sigma \\nu} \\Upsilon^\\sigma_{\\;\\; \\mu \\alpha} + \\frac12 g_{\\mu\\nu}\\Upsilon^{\\alpha \\beta \\sigma} \\Upsilon_{\\alpha \\sigma \\beta} \\nonumber \\end{eqnarray}  This expression is not manifestly symmetric in $\\mu$ and $\\nu$, but the left-hand side is symmetric in its lower indices and hence the right-hand side must be as well.  Thus we use a symmetrized expression to ensure this.  Define for general $K_{\\mu\\nu}$, the symmetrized tensor by $K_{(\\mu\\nu)}=\\frac12(K_{\\mu\\nu}+K_{\\nu\\mu})$.   Using (8) we see that the field equations may be also expressed in the form \\begin{equation}G_{\\mu\\nu} = \\Upsilon^{\\;\\;\\alpha}_{(\\mu \\;\\; \\nu) ;\\alpha}+  \\Upsilon^\\alpha_{\\;\\;\\sigma (\\nu} \\Upsilon^\\sigma_{\\;\\; \\mu) \\alpha}  + \\frac12 g_{\\mu\\nu}\\Upsilon^{\\alpha \\beta \\sigma} \\Upsilon_{\\alpha \\sigma \\beta}\\quad \\equiv 8\\pi \\bigl(T_{\\rm f}\\bigr)_{\\mu\\nu}  \\end{equation} with free field stress energy tensor $\\mathbf{T}_{\\rm f}$. The terms of $\\mathbf{T}_{\\rm f}$ suggest that, when interpreted in Riemannian geometry, this new geometry produces a stress energy tensor with additional terms that could be the stress energy tensor for dark matter or dark energy [6].   In the presence of sources the Lagrangian is of the form \\begin{equation} {\\cal L}= {\\cal L}_{\\rm f}+{\\cal L}_{\\rm s} =  \\int \\biggl( \\frac1{16\\pi}C^\\alpha C_\\alpha + L_s \\biggr) \\, h \\; d^4 x  \\end{equation}  where $L_s = L_s(x^\\alpha)$ is the appropriate Lagrangian density function for the source. In this case $C_\\alpha$ is nonzero  and variation of (10) with respect to the tetrad results in $$\\int \\Biggl[\\frac{1}{16\\pi}\\biggl(C_{(\\mu ;\\nu)}- C_\\alpha \\Upsilon^\\alpha_{\\; (\\mu\\nu)} -\\frac12 g_{\\mu\\nu}C^\\alpha C_\\alpha - g_{\\mu\\nu}C^\\alpha_{\\; ;\\alpha} \\biggr) -\\frac12 (T_{\\rm s})_{\\mu\\nu}\\Biggr] h\\, h^{i\\nu} \\delta h_i^{\\;\\mu} \\, d^4 x = 0$$ Here, $(T_{\\rm s})_{\\mu\\nu}$ is the usual stress-energy tensor of the source for the standard theory [11].  Thus \\begin{equation} C_{(\\mu ;\\nu)}- C_\\alpha \\Upsilon^\\alpha_{\\; (\\mu\\nu)} -\\frac12 g_{\\mu\\nu}C^\\alpha C_\\alpha - g_{\\mu\\nu}C^\\alpha_{\\; ;\\alpha} = 8\\pi (T_{\\rm s})_{\\mu\\nu} \\end{equation} and also we have the following identity for the Einstein tensor, \\begin{equation} G_{\\mu\\nu}=  \\biggl( \\Upsilon^{\\;\\;\\alpha}_{(\\mu \\;\\; \\nu ) ;\\alpha}+\\Upsilon^\\alpha_{\\;\\;\\sigma (\\nu} \\Upsilon^\\sigma_{\\;\\; \\mu ) \\alpha} + \\frac12 g_{\\mu\\nu}\\Upsilon^{\\alpha \\beta \\sigma} \\Upsilon_{\\alpha \\sigma \\beta} \\biggr) + 8\\pi (T_{\\rm s})_{\\mu\\nu}  \\quad   \\end{equation} or \\begin{equation} G_{\\mu\\nu} =  8\\pi\\bigl(T_{\\rm f}\\bigr)_{\\mu\\nu}  + 8\\pi\\bigl(T_{\\rm s}\\bigr)_{\\mu\\nu}  \\qquad .  \\end{equation}  We call $\\mathbf{T}_{\\rm f}$ the free field stress energy and $\\mathbf{T}_{\\rm s}$ the stress energy for the source.      \\par \\phantom{D} \\par ",
        "conclusions": "The theory based on the conservative transformation group may provide a theoretical basis for a unified field theory and may also provide a theoretical basis for dark matter and the correct modification of general relativity.  The Lagrangian for the field with sources may be used in a variety of applications, including quantization.  The internal solution and its corresponding stellar model needs additional work to produce more realistic models.  The external solutions, being non-compact, show promise for explaining dark matter.  Excellent agreement is found with Kepler's Law and redshift.  The theory also gives a realistic explanation for the Pioneer anomaly and the high temperature of the corona.  While there are differences in the precession of perihelia, light deflection and time delay predictions, these may be explained by the fact that the stress-energy tensor is non-zero, yielding densities and pressures that affect the motion of planets and photons.   \\subsection*"
    },
    "1107/1107.0436_arXiv.txt": {
        "abstract": "High-time-resolution X-ray observations of compact objects provide direct access to strong-field gravity, to the equation of state of ultradense matter and to black hole masses and spins. A 10 m$^{2}$-class instrument in combination with good spectral resolution is required to exploit the relevant diagnostics and answer two of the fundamental questions of the European Space Agency (ESA) Cosmic Vision Theme ``Matter under extreme conditions'', namely: does matter orbiting close to the event horizon follow the predictions of general relativity? What is the equation of state of matter in neutron stars? The Large Observatory For X-ray Timing (LOFT), selected by ESA as one of the four Cosmic Vision M3 candidate missions to undergo an assessment phase, will revolutionise the study of collapsed objects in our galaxy and of the brightest supermassive black holes in active galactic nuclei. Thanks to an innovative design and the development of large-area monolithic silicon drift detectors, the Large Area Detector (LAD) on board LOFT will achieve an effective area of $\\sim$12 m$^{2}$ (more than an order of magnitude larger than any spaceborne predecessor) in the 2--30~keV range (up to 50~keV in expanded mode), yet still fits a conventional platform and small/medium-class launcher. With this large area and a spectral resolution of $<$260 eV, LOFT will yield unprecedented information on strongly curved spacetimes and matter under extreme conditions of pressure and magnetic field strength. ",
        "introduction": "\\label{s:introduction} Our knowledge of compact objects derives from the discoveries made with a steady progression of ever larger-area, hence more sensitive, X-ray timing instrumentation, from Uhuru (US, 0.08 m$^2$, \\cite{Giacconi1971}), via satellites such as EXOSAT (ESA, \\cite{Peacock1981}) and Ginga (Japan, \\cite{Turner1989}) to, most recently, RXTE (US, 0.6 m$^2$, \\cite{Swank2006}). The latter instrument, thanks to it having the largest area so far, finally made good on the promise (e.g. \\cite{Shvartsman1971}) that the process of accretion onto stellar mass collapsed objects should exhibit variability at the millisecond dynamical time scale of the small (few Schwarzschild radii, i.e., kilometres) regions where most of the gravitational potential energy is released. Quasi-periodic oscillations (QPOs) at this time scale were found in the X-ray flux of both neutron stars (kHz QPOs, up to 1250~Hz) and black holes (high frequency QPOs, up to 450~Hz) indicating the action of mechanisms in the strong field region picking out specific frequencies, e.g., orbital and epicyclic motion at specific radii near the marginally stable orbit ($r_\\mathrm{isco}= 6\\,GM/c^2$ for a Schwarzschild black hole) as predicted in General Relativity. The long-sought millisecond spins in accreting neutron stars were also finally found, both during thermonuclear bursts (burst oscillations) and in the persistent flux (coherent periodic signals of millisecond pulsars), as were the vibrational responses to cataclysmic events of some neutron stars themselves (magnetar crust/core oscillations). Relativistically broadened Fe K-lines observed from the accretion disks of both X-ray binaries and active galactic nuclei (AGNs) are thought to arise in the same inner regions of the accretion flow. All these phenomena are direct diagnostics of the motion of matter and the propagation of radiation in very strong gravitational fields: they encode information about the mass, radius and spin of collapsed objects and can probe the properties of ultradense matter and superstrong magnetic fields in neutron stars. Yet tantalizingly, while the relevant signals were identified, due to its still limited sensitivity, RXTE was only able to scratch the surface when it came to actually using these diagnostics to probe strong field gravity and ultradense matter.  LOFT is specifically designed to exploit the diagnostics of very rapid X-ray flux and spectral variability that directly probe the motion of matter down to distances very close to black holes and neutron stars. Its $\\sim$20 times larger effective area than RXTE's PCA  is crucial in this respect. The key to this breakthrough in effective area resides in the synergy between technologies imported from other fields of scientific research. The crucial characteristic of the LOFT Large Area Detector (LAD) is a mass per unit surface in the range of $\\sim$10 kg~m$^{-2}$, enabling a payload with $\\sim$15~m$^{2}$ geometric area at reasonable weight. The ingredients for a sensitive but lightweight experiment are the large-area Silicon Drift Detectors (SDDs) designed on the heritage of the ALICE experiment at CERN/LHC \\cite{Vacchi1991}, and a collimator based on lead-glass capillary plates. The drift concept makes the spectroscopic performance of the SDDs weakly dependent on the extent of the collecting surface: large-area ($\\sim$70~cm$^{2}$) monolithic detectors can be designed, with only 256 read-out anodes (thus low power requirements), but very good spectral performance (FWHM $\\sim$260~eV) \\cite{Zampa2011}.  An unprecedentedly large throughput ($\\sim3 \\times 10^{5}$ cts~s$^{-1}$ from the Crab) is achieved with a segmented detector, making pile-up and dead-time, often worrying or limiting focused experiments, secondary issues. The large detector on LOFT is deployed in space through a mechanism relying on the experience in Synthetic Aperture Radar missions (e.g., the SMOS ESA mission, in orbit since 2009, \\cite{Bueno2005}), where very large panels are deployed in space with high accuracy. This mechanism, combined with the low-weight detector technology, allows the stowing of the LOFT satellite inside the fairing of a Vega launcher to design a small mission in a low-background equatorial LEO ($\\sim$600~km). The LOFT scientific payload is completed by a coded-mask Wide Field Monitor (WFM), that uses the same type of detectors, and is in charge of monitoring a large fraction of the sky potentially accessible to LAD, to provide the history and context for the sources observed by LAD and to trigger its pointed observations on their most interesting and extreme states. Its sensitivity and large sky coverage make the WFM also an important resource on its own right.  In this paper we describe the LOFT scientific objectives and expected scientific performance (Sect.~\\ref{s:scientific_objectives}), mission profile (Sect.~\\ref{s:profile}) and details of the instrumentation and model payload designed to achieve the science objectives (Sect.\\ref{s:model_payload}). ",
        "conclusions": "The LOFT mission proposal was selected by ESA as one of the four candidate M3 missions for the Cosmic Vision programme. ESA and the LOFT consortium will carry out an assessment phase of the mission over $\\sim$18 months, concluding by the end of 2012. Further down-selections of the mission candidates are currently expected to bring only one of the four M3 candidates to a launch in the time-frame 2020-2022.  LOFT was primarily designed to address the fundamental questions raised by the Cosmic Vision under the \"matter under extreme conditions\" theme, through the study of the spectral and fast flux variability of compact X-ray sources. The high throughput required by this type of investigation is achieved by an unprecedented large collecting area: 15 m$^{2}$ geometric area, leading to 12~m$^{2}$ effective area peak at 6-10 keV (see Fig.~\\ref{f:LAD_eff_area}), 20 times larger than any predecessors. The enabling technology is provided by the large area Silicon Drift Detectors and the capillary plate X-ray collimators, offering a large collecting and collimated area with limited resources in terms of mass, volume, power and costs. This type of detectors combine large area with very good spectral performance, enabling spectral studies with a resolution of $\\sim$250~eV (FWHM). The segmentation of the Si drift detectors is a straightforward solution to the issues of dead time and pile-up, among the main problems in timing experiments.  Time variability from sub-ms to days will be explored by the LOFT LAD and WFM with statistical accuracy never achieved to date. For example, for the first time quasi-periodic oscillations in neutron star or black-hole systems will be studied in the time domain (as opposite to the frequency domain), within their coherence timescale. The large effective area and the spectral capabilities of LOFT will make optimal use of the diagnostic potential of the fast variability of the X-ray emission to probe the behaviour of matter in the strong gravity and ultradense environments, such as the innermost stable orbits around black holes or the interior of neutron stars, opening access to the understanding of the physical laws at work in such extreme conditions."
    },
    "1107/1107.5814_arXiv.txt": {
        "abstract": "The atmospheric parameters of the components of the 16\\,Cygni binary system, in which the secondary has a gas giant planet detected, are measured accurately using high quality observational data. Abundances relative to solar are obtained for 25 elements with a mean error of $\\sigma(\\mathrm{[X/H]})=0.023$\\,dex. The fact that \\cyga\\ has about four times more lithium than \\cygb\\ is normal considering the slightly different masses of the stars. The abundance patterns of \\cyga\\ and B, relative to iron, are typical of that observed in most of the so-called solar twin stars, with the exception of the heavy elements $(Z>30)$, which can, however, be explained by Galactic chemical evolution. Differential (A--B) abundances are measured with even higher precision ($\\sigma(\\Delta\\mathrm{[X/H]})=0.018$\\,dex, on average). We find that \\cyga\\ is more metal-rich than \\cygb\\ by $\\Delta\\mathrm{[M/H]}=+0.041\\pm0.007$\\,dex. On an element-to-element basis, no correlation between the A--B abundance differences and dust condensation temperature ($\\tc$) is detected. Based on these results, we conclude that if the process of planet formation around \\cygb\\ is responsible for the observed abundance pattern, the formation of gas giants produces a constant downwards shift in the photospheric abundance of metals, without a $\\tc$ correlation. The latter would be produced by the formation of terrestrial planets instead, as suggested by other recent works on precise elemental abundances. Nevertheless, a scenario consistent with these observations requires the convective envelopes of $\\simeq1M_\\odot$ stars to reach their present-day sizes about three times quicker than predicted by standard stellar evolution models. ",
        "introduction": "It is often assumed that the members of multiple stellar systems with similar components have the exact same chemical composition since they formed from a common molecular cloud and have undergone similar evolutionary paths. A number of studies have found, however, that elemental abundance differences in binary stars are not uncommon \\citep[e.g.,][]{gratton01,laws01,sadakane03,desidera04,desidera06}. Systematic errors in the spectroscopic analysis could be blamed for some of these anomalies but small differences have been observed even in pairs with very similar components, i.e., in cases where differential analysis is expected to minimize the errors. The origin of these abundance anomalies remains unknown.  Stars and their planets form nearly at the same time. Therefore, it is likely that the process of planet formation leaves detectable signatures in the chemical composition of the host stars. The well-known planet-metallicity correlation in late-type main-sequence stars, for example, suggests that high metallicity increases the probability of forming planets, at least the gas giants \\citep[e.g.,][]{marcy05,udry07,johnson10}. \\citet[][hereafter M09]{melendez09:twins} have further suggested that small anomalies detected in the solar chemical composition compared to that of the majority of so-called solar twin stars, objects with spectra nearly indistinguishable from the solar spectrum, could be due to the formation of terrestrial planets \\citep[see also][hereafter R09]{ramirez09}. They detect a small deficiency of refractory elements in the solar photosphere which would disappear if the rocky material that once formed in the solar system were diluted into the present-day solar convective envelope \\citep{chambers10}. Thus, the missing mass of refractories in the Sun seems to have been locked-up in the rocky bodies of the solar system, including the terrestrial planets.  The hypothesis described above assumes that the formation of the massive gas giants in the solar system did not affect the solar chemical composition, which seems unlikely given the much larger masses of these planets in comparison with the terrestrial planets. Nevertheless, this could be the case if, for example, the abundance ratios of different metals in the gas giants combined were similar to those found in the Sun. Unfortunately, the detailed chemical composition of the gas giants, in particular their cores, is not yet known with enough precision to rule out or confirm this assumption. Careful spectroscopic studies of stars with gas giant planets detected could therefore provide us with important clues about the impact of planet formation on the chemical composition of the host stars.  The binary star 16\\,Cygni is an ideal system to look for small elemental abundance anomalies and to study the impact of gas giant planet formation. 16\\,Cyg has long been known as the nearest bound pair of solar twin stars \\cite[e.g.,][]{friel93,porto97,hauser99}, although strictly speaking they should be referred to as solar analogs \\cite[cf.][]{melendez07:twins}, in particular because they appear to be older than the Sun. Careful line-by-line differential analysis of this type of stars yields elemental abundances with the highest precision possible, as demonstrated by the works of M09 and R09. Moreover, the components of the 16\\,Cyg system are bright ($V\\sim6$\\,mag) and therefore very high quality spectra can be obtained for them using mid-sized telescopes. Thus, extremely high precision differential spectroscopic analysis can be applied to this system.  The secondary star in the 16\\,Cyg binary system, \\cygb, is known to host a gas giant planet with a minimum mass of 1.5 Jupiter masses (1.5\\,$M_\\mathrm{J}$) in a highly eccentric ($e\\simeq0.6$) orbit. \\cite{cochran97} were the first to detect this planet and they also obtained accurate radial velocity data that showed that \\cyga\\ does not host the same type of planet (gas giant in short-period orbit). Up-to-date follow-up observations of this pair of stars have not changed their planetary status (W.~Cochran, private communication). The presence of low-mass and/or long orbital period planets around either of these two stars is, of course, not to be ruled out.  A faint companion, likely a cool M-dwarf, to the primary star \\cyga, at an angular distance of about 3.4\\,arcsec has been reported by a number of authors \\citep[e.g.,][]{hauser99,turner01,patience02}. The absolute magnitude difference between \\cyga\\ and C (the faint companion) is about 8.6 mag in the $R$ band. Their reported separation is large enough for the light from 16\\,Cyg\\,C to not interfere with that from \\cyga\\ in our data but even if 16\\,Cyg\\,C was superposed with \\cyga\\ at the time of our observation, it would affect the combined spectrum by less than 0.04\\,\\% at $\\lambda=6000$\\,\\AA\\ and even less at shorter wavelengths.  The atmospheric parameters of the 16\\,Cyg pair have been measured by a number of authors using a variety of techniques. Few of these previous studies, however, present results based on extremely high precision analysis, i.e., significantly better than ``standard'' methods. One of these studies is that by \\cite{laws01}, who find that \\cyga\\ is more iron-rich than \\cygb\\ by about $6\\pm2$\\,\\%. They also comment as a note added in proof that a significant positive correlation with dust condensation temperature is found for the abundance difference of 13 elements analyzed. They claim that their results support a self-pollution model in which \\cyga\\ accreted planetary material that once formed there. They also argue that this scenario could explain the observed difference in lithium abundance between \\cyga\\ and B \\cite[e.g.,][]{king97}. Later work by \\cite{takeda05}, however, suggests that \\cyga\\ and B have exactly the same chemical composition and attributed the discrepancy with the results by \\citeauthor{laws01} to differences in the adopted stellar parameters.  Here we derive atmospheric parameters and elemental abundances as precisely and accurately as possible using high quality data of \\cyga\\ and B. High precision elemental abundances are presented for 25 elements in the 16\\,Cyg binary system. We discuss scenarios that could explain the observed abundance patterns. ",
        "conclusions": ""
    },
    "1107/1107.2119_arXiv.txt": {
        "abstract": "The environs of supermassive black holes are among the universe's most extreme phenomena. Understanding the physical processes occurring in the vicinity of black holes may provide the key to answer a number of fundamental astrophysical questions including the detectability of strong gravity effects, the formation and propagation of relativistic jets, the origin of the highest energy gamma-rays and cosmic-rays, and the nature and evolution of the central engine in Active Galactic Nuclei (AGN). As a step towards this direction, this paper reviews some of the progress achieved in the field based on observations in the very high energy domain. It particularly focuses on non-thermal particle acceleration and emission processes that may occur in the rotating magnetospheres originating from accreting, supermassive black hole systems. Topics covered include direct electric field acceleration in the black hole's magnetosphere, ultra-high energy cosmic ray production, Blandford-Znajek mechanism, centrifugal acceleration and magnetic reconnection, along with the relevant efficiency constraints imposed by interactions with matter, radiation and fields. By way of application, a detailed discussion of well-known sources (Sgr~A*; Cen~A; M87; NGC1399) is presented. ",
        "introduction": "% This review focuses on non-thermal particle acceleration and very high energy (VHE) emission processes that may occur in the vicinity of magnetized supermassive black holes. As such, active galaxies and, to some extent, extinct or dormant quasars are placed in the center of its attention. Although rotating, supermassive black holes have often been considered as putative sources for the energization of ultra-high energy (UHE) cosmic rays, non-thermal magnetospheric models have recently gained an additional impetus with the detection of very high energy gamma-rays from the (non-blazar) radio galaxy M87, located $\\sim 16$ Mpc away in the Virgo cluster of galaxies.\\cite{aharonian06,acciari08,albert08} While former observations of rapid VHE flux variations (on timescales of $\\sim [1-2]$ days) already indicated a small size of the $\\gamma$-ray emitting region, the location of this region remained ambiguous based on the VHE results alone. Additional high-resolution ($\\sim 50$ Schwarzschild radii) radio VLBA imaging has now shown that the gamma-ray flaring activity is accompanied by an increase in the radio flux close to core (see Fig.~\\ref{m87_lc}), indicating that the required energetic charged particles may in fact be accelerated in the very vicinity of the central black hole.\\cite{acciari09} At present, further evidence is needed to strengthen this inference. Nevertheless, magnetospheric models where the relevant non-thermal processes are considered to occur at the base of a rotating black-hole-jet magnetosphere (e.g., refs.\\cite{neronov07,rieger08a,lev11}), have emerged as attractive candidates. If verified by further observations, VHE gamma-ray observations could provide a fundamental diagnostic of the most violent region in Active Galactic Nuclei (AGN).\\\\ \\begin{figure}[t] \\center \\psfig{figure=fig1.pdf, width=10truecm} \\caption{Combined VHE gamma-ray, X-ray and radio light curves for the radio galaxy M87 in 2007--2008. A VHE gamma-ray flare (Fig. A) is followed by an increase of the radio flux (Fig. C) close to the black hole. Knot HST-1, located in the jet $\\sim100$ pc away, is in a low X-ray state (Fig. B) and therefore unlikely to be the source of the observed VHE gamma-rays. Details are as follows: (A) VHE -ray flux data ($E > 0.35$ TeV, nightly average) including H.E.S.S., MAGIC, and VERITAS. The inlay shows a zoomed version of the rapid flaring activity (with timescales as low as $\\sim [1-2]$ days) in February 2008. (B) X-ray (Chandra) flux for the nucleus and the knot HST-1. (C) Radio (43 GHz VLBA) flux data. The shaded horizontal area indicates the range of radio fluxes from the nucleus before the 2008 flare. The radio flux of the outer jet regions does not change substantially; most of the observed increase results from the region around the nucleus. Figure adapted from Acciari et al.\\protect\\cite{acciari09}} \\label{m87_lc} \\end{figure} This review is structured as follows: Sect.~2 gives a phenomenologically-orientated introduction into the close environment of supermassive black holes. In Sect.~3, particle acceleration scenarios that draw on rotating, large-scale poloidal magnetic fields (i.e., the magnetosphere) are presented. Sect.~4 discusses possible interactions of accelerated particles with ambient matter, radiation and fields, that are expected to limit achievable particle energies. Applications to concrete sources as potential VHE gamma-ray and UHE cosmic ray emitters are discussed in Sect.~5. ",
        "conclusions": "Non-thermal processes occurring in the vicinity of supermassive black holes may allow new insights into the behavior of matter and energy under extreme conditions. The detection of, e.g., rapid variability on timescales of  $\\sim r_s/c$ and unusual high-energy emission signatures can provide an important diagnostic tool towards this end. Based on this, VHE observations during the last years have demonstrated the capacities of the new generation of VHE gamma-ray instruments (e.g., MAGIC II, HESS II, VERITAS and FERMI) to substantially improve our knowledge of the activity center in nearby, underluminous AGNs and to trigger new theoretical developments that could serve as benchmark for future VHE arrays like CTA. Detailed theoretical modelling of non-thermal particle acceleration and radiative numerical simulations may soon become indispensable in order to fully take advantage of the improved observational capacities."
    },
    "1107/1107.2269_arXiv.txt": {
        "abstract": "We present the first observations of emission lines of CN(2-1), HCO$^{+}$(3-2) and C$_{2}$H(3-2) in the Perseus cluster. We observed at two positions: directly at the central galaxy, NGC 1275 and also at a position about 20$''$ to the east where associated filamentary structure has been shown to have strong CO emission. Clear detections in CN and HCO$^{+}$ transitions and a weak detection of the C$_{2}$H transition were made towards NGC 1275, while weak detections of CN and HCO$^{+}$ were made towards the eastern filamentary structure. Crude estimates of the column densities and fractional abundances (mostly upper limits) as functions of an unknown rotational temperature were made to both sources. These observational data were compared with the outputs of thermal/chemical models previously published by \\citet{Baye10c} in an attempt to constrain the heating mechanisms in cluster gas. We find that models in which heating is dominated by cosmic rays can account for the molecular observations. This conclusion is consistent with that of \\citet{Ferl09} in their study of gas traced by optical and infrared radiation. The cosmic ray heating rate in the regions probed by molecular emissions is required to be at least two orders of magnitude larger than that in the Milky Way. ",
        "introduction": "\\label{sec:intro}   The Perseus Cluster is one of the nearest galaxy clusters and is the brightest X-ray cluster in the sky. The cluster and its central galaxy NGC 1275 have been the focus of intense study for many years, at X-ray, optical, IR and millimetre wavelengths. The first molecular detections were of CO rotational emission towards the centre of NGC 1275 \\citep{Laza89, Mira89, Reut93, Brai95, Inou96, Brid98, Lim08}. Observations have now demonstrated that CO emission also extends to some tens of kpc from the central galaxy \\citep{Salo06,Salo08a,Salo08b,Salo11} and is strongly correlated with the filamentary structure observed in H$\\alpha$ \\citep{Hu83, Cowi83,Cons01} within the hot gas detected in X-rays at 0.5 keV \\citep{Fabi08}. The H$\\alpha$ structure is also correlated with warm H$_{2}$ emission \\citep{Edge02,Wilm02}. Some of the molecular gas towards the cluster must be at high density ($\\geqslant 10^{4}$ cm$^{-3}$). \\citet{Salo08a} have made the first detection of the high density tracer HCN(3-2) emission towards the centre of NGC 1275.   The nature of the filamentary structure is the subject of current discussion. \\citet{Salo11} note two possibilities: either that the CO filaments form far from the galaxy's centre from uplifted warm gas, eventually falling back \\citep{Reva08}, or that the molecular gas is entrained and dragged out of the galaxy by rising hot gas. Evidently, further observations of molecular gas may help to determine its origin. In this paper, we present observations of emission lines of CN(2-1), HCO$^{+}$(3-2) and C$_{2}$H(3-2) towards the central galaxy and also towards a position 20$''$ to the east where there is strong CO emission in the associated filamentary structure.   These three species were identified in a theoretical study \\citep{Baye10c} as tracers of regions influenced by the dissipation of turbulence and waves, heating the gas and accelerating cosmic rays (see also \\citealt{Craw92, Pope08a}). \\citet{Ferl09} developed models of the optical and infrared emission filamentary regions in order to infer the rate at which mechanical energy is dissipated, and to determine the cosmic ray background produced by the interaction. \\citet{Ferl09} concluded that the heating rate per hydrogen nucleus due to dissipation or cosmic ray induced ionisation must be 300 times higher in the optical emission filaments than in the local interstellar medium.   In their work, \\citep{Baye10c} identified the species CN, C$_{2}$H and HCO$^{+}$ as sensitive to changes in cosmic ray ionisation rate or to an additional source of heating such as dissipation. \\citet{Meij11} came to a similar conclusion. Here, we present the observational follow-up to that theoretical study; we aim to determine whether dissipation or cosmic ray induced ionisation dominates the heating.   The paper is organised as follows: in Section \\ref{sec:obs} we present the first observations of CN(2-1), C$_{2}$H(3-2) and HCO$^{+}$(3-2) emissions in both NGC 1275 and a position 20$''$ to the east where emission in the CO filamentary structure is known to be strong (see Figure \\ref{fig:0}). In Section \\ref{sec:resu} we analyse the data and provide rough estimates of column densities and fractional abundances of the three molecules in NGC 1275 and in the filamentary region. Section \\ref{sec:comp} compares the results of the observations with the models of \\citet{Baye10c} to try to identify whether cosmic ray heating or dissipation is the dominant heating mechanism, or whether a combination of both is required. Section \\ref{sec:con} discusses the results and gives our conclusions. ",
        "conclusions": "\\label{sec:con}   The comparison in Section \\ref{sec:comp} between the molecular abundances obtained from the observations and from the models of \\citet{Baye10c} suggests that we can identify one model, Model E, that produces results consistent with the observations. This model has a low FUV radiation field, a low rate of heating by dissipation, but a high cosmic ray ionization rate. Unpublished data from the \\citet{Baye10c} calculations show that cosmic ray heating accounts for 85\\% of all heating at either $A_{v}=$ 3 or 8 mag.   The enhanced cosmic ray flux may arise in a dynamical interaction, as suggested by \\citet{Ferl09}. Those authors also found that an enhanced cosmic ray ionization rate was required to account for the observations of the optical and infrared emitting components of the filaments. It is interesting, therefore, that the results from the present work also suggest that the regions of cluster gas probed by millimetre and sub-millimetre emissions also require a similarly enhanced cosmic ray ionization rate. A second inference is that a high heating rate from energy dissipation is not required; in fact, as Table \\ref{tab:sum} indicates, a high heating rate from sources other than cosmic rays appears to inhibit a match between models and observations.   In the observations, we have detected only one line in each of three molecular species, so a reliable analysis of the observational data cannot be made. It would be useful for the further study of galaxy cluster gas to make confirmed detections of several lines of each of the three molecular tracers used in this work. In the modelling of the chemistry of the gas in and around NGC 1275 (\\citealt{Baye10c}), some important physical parameters (e.g. gas:dust ratio and metallicity) are poorly known and undoubtedly have an impact on the model predictions. Nevertheless, the present work shows the value of molecular line observations of cluster gas, and it is interesting that the results reported here for the molecular gas are reasonably consistent with the results for regions probed by the optical and infrared. The main conclusions of this work are that a cosmic ray ionization rate enhanced by a factor of at least one hundred in the cluster gas is required to establish the observed molecular abundances, and that a high heating rate from other causes appears to mitigate against that chemistry. In these circumstances, the HCO$^{+}$ emission is strong.   \\begin{table*} \\caption{Total column density estimates (in cm$^{-2}$) obtained using the formula described in Section \\ref{sec:resu}, for various values of $T_{rot}$. Since the there are observed upper limits for both NGC 1275 and the filamentary region, the corresponding derived total column densities presented here are thus also to be considered as upper limits (symbol $<$ used). The values of $T_{rot}$ has been selected with respect to the available online values of the partition function $Q(T_{rot})$ and such as to be compatible with the range of kinetic temperature displayed by the models (see Table \\ref{tab:mod}).}\\label{tab:resu} \\begin{tabular}{l c c c c c c c} \\hline & $T_{rot}=2.725$ K & $T_{rot}= 5$ K & $T_{rot}= 9.375$ K & $T_{rot}= 18.75$ K & $T_{rot}= 37.5$ K & $T_{rot}= 50$ K & $T_{rot}= 75$ K\\\\ \\hline NGC 1275 &&&&&&&\\\\ \\hline HCO$^{+}$ & $6.93 \\times 10^{14}$ & $1.84 \\times 10^{13}$ & $2.62 \\times 10^{12}$ & $9.45 \\times 10^{11}$ & $6.22 \\times 10^{11}$ & $5.86 \\times 10^{11}$ & $5.46 \\times 10^{11}$\\\\ CN & $2.42 \\times 10^{14}$ & $1.91 \\times 10^{13}$ & $5.02 \\times 10^{12}$ & $2.53 \\times 10^{12}$ & $1.93 \\times 10^{12}$ & $1.86 \\times 10^{12}$ & $1.79 \\times 10^{12}$\\\\ C$_{2}$H & $<9.03 \\times 10^{15}$ & $<1.55 \\times 10^{14}$ & $<1.74 \\times 10^{13}$ & $<5.40 \\times 10^{12}$ & $<3.23 \\times 10^{12}$ & $<2.94 \\times 10^{12}$ & $<2.66 \\times 10^{12}$\\\\ \\hline \\hline Filament &&&&&&&\\\\ \\hline HCO$^{+}$ & $<2.10 \\times 10^{14}$ & $<5.58 \\times 10^{12}$ & $<7.94 \\times 10^{11}$ & $<2.87 \\times 10^{11}$ & $<1.88 \\times 10^{11}$ & $<1.78 \\times 10^{11}$ & $<1.66 \\times 10^{11}$\\\\ CN & $<4.68 \\times 10^{14}$ & $<3.69 \\times 10^{13}$ & $<9.70 \\times 10^{12}$ & $<4.89 \\times 10^{12}$ & $<3.73 \\times 10^{12}$ & $<3.60 \\times 10^{12}$ & $<3.36 \\times 10^{12}$\\\\ C$_{2}$H & $<1.34 \\times 10^{16}$ & $<2.52 \\times 10^{14}$ & $<2.97 \\times 10^{13}$ & $<9.47 \\times 10^{12}$ & $<5.74 \\times 10^{12}$ & $<5.25 \\times 10^{12}$ & $<4.77 \\times 10^{12}$\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*} \\caption{Fractional abundances (with respect to the total number of hydrogen nuclei) obtained using hypothesis described in Section \\ref{sec:resu}, for various values of $T_{rot}$ and for $A_{v}=3$ mag and 8 mag (consistently with model predictions from \\citealt{Baye10c}. See text in Section \\ref{sec:comp}). See caption of Table \\ref{sec:resu}.}\\label{tab:abun} \\begin{tabular}{l c c c c c c c} \\hline & $T_{rot}=2.725$ K & $T_{rot}= 5$ K & $T_{rot}= 9.375$ K & $T_{rot}= 18.75$ K & $T_{rot}= 37.5$ K & $T_{rot}= 50$ K & $T_{rot}= 75$ K\\\\ \\hline NGC 1275 &&&&&&&\\\\ \\hline $A_{v}=3$ mag &&&&&&&\\\\ HCO$^{+}$ & $1.44 \\times 10^{-7}$ & $3.84 \\times 10^{-9}$ & $5.45 \\times 10^{-10}$ & $1.97 \\times 10^{-10}$ & $1.30 \\times 10^{-10}$ & $1.22 \\times 10^{-10}$ & $1.14 \\times 10^{-10}$\\\\ CN & $5.04 \\times 10^{-8}$ & $3.97 \\times 10^{-9}$ & $1.05 \\times 10^{-9}$ & $5.27 \\times 10^{-10}$ & $4.01 \\times 10^{-10}$ & $3.87 \\times 10^{-10}$ & $3.72 \\times 10^{-10}$\\\\ C$_{2}$H & $<1.88 \\times 10^{-6}$ & $<3.23 \\times 10^{-8}$ & $<3.62 \\times 10^{-9}$ & $<1.12 \\times 10^{-9}$ & $<6.72 \\times 10^{-10}$ & $<6.12 \\times 10^{-10}$ & $<5.54 \\times 10^{-10}$\\\\ \\hline $A_{v}=8$ mag &&&&&&&\\\\ HCO$^{+}$ & $5.42 \\times 10^{-8}$ & $1.44 \\times 10^{-9}$ & $2.05 \\times 10^{-10}$ & $7.39 \\times 10^{-11}$ & $4.86 \\times 10^{-11}$ & $4.58 \\times 10^{-11}$ & $4.27 \\times 10^{-11}$\\\\ CN & $1.89 \\times 10^{-8}$ & $1.49 \\times 10^{-9}$ & $3.92 \\times 10^{-10}$ & $1.98 \\times 10^{-10}$ & $1.51 \\times 10^{-10}$ & $1.45 \\times 10^{-10}$ & $1.40 \\times 10^{-10}$\\\\ C$_{2}$H & $<7.06 \\times 10^{-7}$ & $<1.21 \\times 10^{-8}$ & $<1.36 \\times 10^{-9}$ & $<4.22 \\times 10^{-10}$ & $<2.52 \\times 10^{-10}$ & $<2.29 \\times 10^{-10}$ & $<2.08 \\times 10^{-10}$\\\\ \\hline \\hline Filament &&&&&&&\\\\ \\hline $A_{v}=3$ mag &&&&&&&\\\\ HCO$^{+}$ & $<4.38 \\times 10^{-8}$ & $<1.16 \\times 10^{-9}$ & $<1.65 \\times 10^{-10}$ & $<5.97 \\times 10^{-11}$ & $<3.93 \\times 10^{-11}$ & $<3.70 \\times 10^{-11}$ & $<3.45 \\times 10^{-11}$\\\\ CN & $<9.74\\times 10^{-8}$ & $<7.68 \\times 10^{-9}$ & $<2.02 \\times 10^{-9}$ & $< 1.02 \\times 10^{-9}$ & $<7.76 \\times 10^{10}$ & $<7.50 \\times 10^{-10}$ & $<7.00 \\times 10^{-10}$\\\\ C$_{2}$H & $<2.79 \\times 10^{-6}$ & $<5.24 \\times 10^{-8}$ & $<6.18 \\times 10^{-9}$ & $<1.97 \\times 10^{-9}$ & $<1.20 \\times 10^{-9}$ & $<1.09 \\times 10^{-9}$ & $<9.93 \\times 10^{-10}$\\\\ \\hline $A_{v}=8$ mag &&&&&&&\\\\ HCO$^{+}$ & $<1.64 \\times 10^{-8}$ & $<4.36 \\times 10^{-10}$ & $<6.20 \\times 10^{-11}$ & $<2.24 \\times 10^{-11}$ & $<1.47 \\times 10^{-11}$ & $<1.39 \\times 10^{-11}$ & $<1.29 \\times 10^{-11}$\\\\ CN & $<3.65 \\times 10^{-8}$ & $< 2.88 \\times 10^{-9}$ & $< 7.58 \\times 10^{-10}$ & $<3.82 \\times 10^{-10}$ & $<2.91 \\times 10^{-10}$ & $<2.81 \\times 10^{-10}$ & $<2.62 \\times 10^{-10}$\\\\ C$_{2}$H & $<1.05 \\times 10^{-6}$ & $<1.97 \\times 10^{-8}$ & $<2.32 \\times 10^{-9}$ & $<7.40 \\times 10^{-10}$ & $<4.49 \\times 10^{-10}$ & $<4.10 \\times 10^{-10}$ & $<3.72 \\times 10^{-10}$\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*} \\caption{Predicted values of the fractional abundances (with respect to the total number of hydrogen nuclei) of CN, C$_{2}$H and HCO$^{+}$ from models seen in \\citet{Baye10c}. More precisely, Models A - E listed here correspond to Models 16, 14, 5, 18, and 12, respectively, in \\citet{Baye10c} hence to models having high values of both $H$ and $\\zeta$ with a low value of $I$; low values of $H$, $\\zeta$ and $I$; low values of $H$ and $\\zeta$ with a high value of $I$; high value of $H$ and low values of $\\zeta$ and $I$ and low value of $H$, high value of $\\zeta$ and low value of $I$, respectively. See the text in Section \\ref{sec:comp}.}\\label{tab:mod} \\begin{tabular}{l c c c c c c c c} \\hline Model & $A_{v}$ & $T_{K}$ & $H$ & $\\zeta$ & $I$ & Frac. abun. & Frac. abun. & Frac. abun.\\\\ & (mag) & (K) & (ergcm$^{-3}$ s$^{-1}$) & (s$^{-1}$) & (Habing) & CN & C$_{2}$H & HCO$^{+}$\\\\ \\hline A & 3 & 43.8 & $1.00 \\times 10^{-20}$ & $5.00 \\times 10^{-15}$ & 10 & $7.33 \\times 10^{-9}$ & $3.32 \\times 10^{-10}$ & $4.69 \\times 10^{-10}$\\\\ A & 8 & 47.5 & $1.00 \\times 10^{-20}$ & $5.00 \\times 10^{-15}$ & 10 & $7.24 \\times 10^{-9}$ & $2.74 \\times 10^{-10}$ & $5.94 \\times 10^{-10}$\\\\ \\hline B & 3 & 13.8 & $1.00 \\times 10^{-22}$ & $5.00 \\times 10^{-17}$ & 10 & $1.74 \\times 10^{-9}$ & $9.01 \\times 10^{-12}$ & $1.36 \\times 10^{-11}$\\\\ B & 8 & 14.9 & $1.00 \\times 10^{-22}$ & $5.00 \\times 10^{-17}$ & 10 & $1.81 \\times 10^{-10}$ & $5.22 \\times 10^{-12}$ & $5.32 \\times 10^{-11}$\\\\ \\hline C & 3 & 25.3 & $1.00 \\times 10^{-22}$ & $5.00 \\times 10^{-17}$ & 100 & $1.55 \\times 10^{-9}$ & $5.50 \\times 10^{-11}$ & $1.31 \\times 10^{-11}$\\\\ C & 8 & 14.5 & $1.00 \\times 10^{-22}$ & $5.00 \\times 10^{-17}$ & 100 & $1.82 \\times 10^{-10}$ & $5.28 \\times 10^{-12}$ & $5.16 \\times 10^{-11}$\\\\ \\hline D & 3 & 74.8 & $1.00 \\times 10^{-20}$ & $5.00 \\times 10^{-17}$ & 10 & $2.51 \\times 10^{-9}$ & $8.82 \\times 10^{-11}$ & $1.65 \\times 10^{-10}$\\\\ D & 8 & 63.3 & $1.00 \\times 10^{-20}$ & $5.00 \\times 10^{-17}$ & 10 & $8.61 \\times 10^{-10}$ & $5.43 \\times 10^{-12}$ & $2.62 \\times 10^{-10}$\\\\ \\hline E & 3 & 51.0 & $1.00 \\times 10^{-22}$ & $5.00 \\times 10^{-15}$ & 10 & $1.82 \\times 10^{-9}$ & $2.68 \\times 10^{-10}$ & $1.37 \\times 10^{-10}$\\\\ E & 8 & 60.2 & $1.00 \\times 10^{-22}$ & $5.00 \\times 10^{-15}$ & 10 & $1.99 \\times 10^{-9}$ & $2.95 \\times 10^{-10}$ & $2.94 \\times 10^{-10}$\\\\ \\hline \\end{tabular} \\end{table*}   \\begin{table*} \\caption{Summary of the conclusions obtained when comparing the estimates of the fractional abundances derived from the observations to the model predictions (see Section \\ref{sec:comp}). We have assumed that models and observations are in agreement (symbol `+' used) when their values differ from less than or equal to a factor of 5 difference. Otherwise, the symbol `-' is used. `(3)' or `(8)' means that the fractional abundances at $A_{v}=3$ mag or at $A_{v}=8$ mag are well reproduced by the models (but not at both $A_{v}$s), respectively.}\\label{tab:sum} \\begin{tabular} {l c c c c c c} \\hline Model  &        C$_{2}$H  & C$_{2}$H & CN & CN & HCO$^{+}$ & HCO$^{+}$\\\\ Parameters & centre & filament & centre & filament & centre & filament\\\\ \\hline Model A: $I$ low, $H$ high, $\\zeta$ high& $+$ & $+$ & -   & -   & (3) & -  \\\\ Model B: $I$ low, $H$ low, $\\zeta$ low  & -   & -   & $+$ & $+$ & (8) & (8)\\\\ Model C: $I$ high, $H$ low, $\\zeta$ low &  -  &  -  & $+$ & $+$ & (8) & (8)\\\\ Model D: $I$ low, $H$ high, $\\zeta$ low & -   & -   & $+$ & $+$ & (3) & (3)\\\\ Model E: $I$ low, $H$ low, $\\zeta$ high & $+$ & $+$ & (3) & (3) & (3) & (3)\\\\ \\hline \\end{tabular} \\end{table*}"
    },
    "1107/1107.5449_arXiv.txt": {
        "abstract": "In this paper we consider the possibility that the structure of the largest radio galaxy \\J is formed by a restarted rather than a primary jet activity. This hypothesis is motivated by the unusual morphological properties of the source, suggesting almost ballistic propagation of powerful jets in a particularly low-density environment. New radio observations of \\J confirm its morphology consisting of only two narrow lobes; no trace of any outer low-density cavity due to the previous jet activity is therefore detected. Different model fits performed using the newly accessed radio data imply relatively young age of the source, its exceptionally high expansion velocity, large jet kinetic power, and confirm particularly low-density environment. We find that it is possible to choose a realistic set of the model parameters for which the hypothetical outer lobes of \\J are old enough so that their expected radio surface brightness is substantially below the rms noise level of the available radio maps. On the other hand, the extremely low density of the gas surrounding the \\J lobes is consistent with the mean density of the baryonic matter in the Universe. This suggests that the source may be instead located in a real void of the galaxy and matter distribution. In both cases the giant radio lobes of \\J are expected to modify substantially the surrounding matter by driving strong shocks and heating the gas located at the outskirts of the filamentary galactic distribution. Finally, we also find that the energetic requirements for the source are severe in terms of the total jet power and the total energy deposited by the outflows far away from the central engine. ",
        "introduction": "Giant radio galaxies (hereafter GRGs) are luminous but low-surface brightness sources dominated at radio wavelengths by the emission of extended lobes with linear sizes $\\gtrsim 1$\\,Mpc. Morphologically, most of GRGs resemble powerful Fanaroff-Riley type II objects \\citep[FR\\,IIs;][]{fan74}, and hence dynamical models developed for classical doubles are typically used to infer their intrinsic parameters based on the multifrequency radio data obtained using a variety of radio instruments \\citep[e.g.,][]{kon08,jam08,mach09}. The key difficulty in modeling and understanding of such sources is however the fact that the main characteristics of the ambient medium surrounding giant lobes are hardly known. In fact, there are almost no observational constraints on the structure, temperature, or density of the intergalactic gas at Mpc distances from the centers of poor clusters or groups in which powerful radio galaxies, including GRGs, are expected to be located \\citep{jam04,sar05}. But one may also turn the argument around, and argue that GRGs can indeed be used \\emph{as unique probes} of the intergalactic medium (IGM) and of its cosmological evolution at the outskirts, or even outside of groups and clusters of galaxies. This idea was discussed and pushed forward by, e.g., \\citet{sub08} and \\citet{saf09}.  Clearly, using GRGs as cosmological probes requires a consensus regarding the nature of those objects. Are therefore GRGs just evolved FR\\,II galaxies, for which large linear sizes are strictly due to their advanced ages of the order of $100$\\,Myr? Or are they instead intrinsically different from classical doubles \\citep[as suggested by, e.g.,][]{kai99}? A detailed analysis of various samples of GRGs favors the former possibility, suggesting however in addition that densities of the medium surrounding the discussed objects are lower than average for the whole FR\\,II population \\citep{mack98,sch00b,mach06}. It is interesting to note in this context an analogy with the sources located at the other extremum of the size distribution of radio galaxies, namely Compact Symmetric Objects (CSOs), which spectroscopically are often classified as GigaHertz-Peaked Spectrum (GPS) sources. For these, the long-debated question was if their linear sizes $0.1 - 10$\\,kpc are due to young ages ($\\leq 0.1$\\,Myr) of otherwise `regular' FR\\,II-like radio structures, or instead are due to exceptionally dense ambient medium frustrating evolved jets and confining their lobes within host galaxies \\citep[see][for a review]{odea98}. Or do they constitute yet another intrinsically distinct population of extragalactic radio sources? The merging agreement is that compact radio galaxies are indeed newly born precursors of classical doubles \\citep[and also of low-power FR\\,I sources; see the discussion in][and references therein]{sta08}, even though some of them may reside in particularly overdense environment \\citep[see, e.g.,][]{gar07}.  During the recent years it became clear that understanding the evolution of extragalactic radio sources --- including the youngest and the oldest examples of this class of astrophysical objects --- involves however not only accessing the properties of the ambient medium into which they evolve, or the main jet parameters (like jet total kinetic power, lifetime, etc.), but also characterizing a complex duty cycle of the jet activity. In particular, population studies of radio galaxies and quasars indicated that the jet activity must be intermittent, or at least highly modulated, in addition on different timescales \\citep[e.g.,][]{rey97,alex00,sch01,cze09}. The direct observational evidences for such an intermittency consist of different manifestations of the interrupted and/or restarted jets in radio-loud Active Galactic Nuclei (AGN), including interestingly some GRGs with compact inner structures of the CSO/GPS type \\citep[e.g.,][]{sch99,kon04,kon08,sai07}. This is an important point indeed, since it indicates that the appearance of large-scale jets and lobes is shaped not exclusively by the jet lifetime and the jet kinetic power relative to the density of the ambient medium, as believed previously, but is a complex convolution of these factors with the jet duty cycle. Note in particular that the restarted jets do not propagate within the undisturbed interstellar/intergalactic medium, but instead within the environment modified by the passage of the outflow during the previous epoch of the jet activity, and this may affect substantially the observed properties of the newly formed lobes.  Double-double radio galaxies (hereafter referred to as DDRGs) are the most striking and clear examples of a highly intermittent activity of the jet engine with no realignment of the outflows' direction in the two subsequent activity epochs \\citep{lar99,sch00a}. These are the sources in which an extended radio structure comprises at least two pairs of lobes with a common core, both strongly co-aligned to within a few degrees to each other. At present, about 15 of such systems have been identified from either radio or X-ray observations \\citep[for a review see][]{sai09}. The outer and also the inner lobes of DDRGs are uniquely characterized by an edge-brightened morphology of the FR\\,II type, what led at the first place to the interpretation of the inner structures in DDRGs as `regular lobes' indeed, inflated by a plasma left behind a head (a termination shock) of a new-born jet which propagates through an outer cocoon formed during the previous epoch of the jet activity \\citep{kai00}. The alternative interpretation for the inner structures of DDRGs would involve ballistic and `naked' (lobeless) jets observed in a few performed 2D and 3D hydrodynamical simulations of restarting outflows in AGN \\citep{cla91,cla97}. Bulk of the observations and analysis carried till now for several best known DDRGs are in general consistent with the former scenario \\citep[see, e.g.,][]{sar02,sar03,kon06,saf08,mach10}, even though some potential problems of the model has been recently discussed in \\citet{bro07,bro11}.  The problems arise when the exact evolution of the inner lobes of DDRGs is analyzed. These inner lobes appear rather slim on the high-dynamics radio maps, displaying in particular axial ratios in the range of $9 - 16$, to be compared with the range of $2 - 5$ established for the outer lobes of same systems, and also for the lobes in `regular' FR\\,II radio galaxies. This indicates that the restarted jets in DDRGs propagate within the environment much rarified with respect to the original IGM in which the outer lobes are evolving, as in fact expected in the framework of the jet intermittency scenario, but yet much denser than the one provided simply by the injection of a jet plasma to the outer lobes in the previous epoch of the jet activity \\citep[see, e.g., the detailed analysis by][]{saf08}. Thus required contamination of the remnant cocoons by the additional matter is one of the central questions regarding the nature of DDRGs \\citep[see the discussion in][]{kai00}. It is not clear, however, if any entrainment process operates on sufficiently short timescales to enrich substantially the outer cocoons with the required amount of the cold matter from the IGM during the lifetime of a radio source. For this reason, \\citet{bro07,bro11} argued that radio emission observed from the inner structures of DDRGs is not due to the lobes of newly born jets, but predominantly due to bow shocks induced within the extended lobes by only briefly interrupted relativistic outflows. In this scenario no contamination of the extended/outer cocoon by the additional matter would be needed.  In this paper we consider and analyze a possibility that the structure of the largest radio galaxy known up to date --- \\J with the total projected linear size of $D=4.69$\\,Mpc, discovered by \\citet{mach08} --- is formed by a restarted, rather than a primary jet activity. The motivation for the performed analysis and new radio observations is due to the fact that the axial ratio for the lobes in this outstanding example of a giant, $R_{\\rm T}\\approx 12$, is of the same order or even larger than axial ratios for the inner lobes in all well studied DDRGs, and therefore very different from the typical axial ratios characterizing `normal size' FR\\,II sources or other known GRGs. Since in the framework of the standard dynamical model for classical doubles the axial ratio of the lobes depends predominantly on the jet kinetic power relative to the density of the surrounding, \\J must be a powerful source evolving in a particularly low-density environment, as indeed concluded in \\citet{mach08}. The question is if the low-density of the matter surrounding giant lobes of \\J is due to a unique location of \\J in a void region, or due to some previous jet activity epoch which took place before the formation of the currently observed lobes, and which rarified the IGM thereby (in analogy with DDRGs). In the former case, \\J may be used as a particularly interesting (due to its enormous size) probe of a `warm-hot phase of the intergalactic medium' (hereafter WHIM) within a filamentary galaxy distribution far away from any group or cluster of galaxies. In the latter case, \\J would be the most extreme example of a restarting radio galaxy, providing interesting constraints on the jet duty cycle and cosmological evolution of radio-loud AGN.  New observations of \\J, conducted with the Giant Metrewave Radio Telescope (GMRT) and Very Large Array (VLA), confirm its morphology consisting of only two narrow lobes discerned in the VLA $1400$\\,MHz sky surveys FIRST \\citep{bec95} and NVSS \\citep{con98}, together with the radio core coincident with a parent galaxy at the redshift of $z=0.3067$. The whole radio structure is highly collinear and symmetric: the arm ratio is 1.08 only, and the misalignment angle is $1.3 \\degr$. No trace of any outer cocoon is however detected in our new observations. We modeled \\J using the \\D algorithm of \\citet{mach07}, along with the three bona-fide DDRGs treated as comparison and control sources. The resulting `radio luminosity---linear size' $(P_{\\nu}-D)$ evolutionary diagrams are used to determine conditions allowing the hypothetical outer lobes in the discussed object to dim in radio fluxes to a level precluding the detection at the given sensitivity limits. The archival data regarding the exemplary DDRGs used in the modeling, as well as the new GMRT and VLA observations of \\J resulting in high-resolution low- and high-frequency maps of the source, are presented in \\S\\,2 of the paper. The applied model and the obtained model results are described in \\S\\,3. Further discussion regarding the jet energetics and the properties of the environment of \\J are given in \\S\\,4. Final conclusions are summarized in \\S\\,5. Distances, linear sizes, volumes, and luminosities of the analyzed sources are calculated for the modern $\\Lambda$CDM cosmology ($\\Omega_{\\rm m}=0.27$, $\\Omega_{\\Lambda}=0.73$, and $H_{0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$). ",
        "conclusions": "In this paper we consider the possibility that the structure of the largest radio galaxy, \\J discovered by \\citet{mach08}, is formed by a restarted, rather than a primary jet activity. This hypothesis is motivated by the unusual morphological properties of the source, suggesting almost ballistic propagation of powerful jets in a particularly low-density ambient medium up to $>2$\\,Mpc distances from the host galaxy. If the low density of the source environment is due to the previous jet activity indeed, or rather due to a unique location of \\J in a void region of the galaxy distribution, is the main question addressed in the analysis. The obtained results and the final conclusions are as follows: \\begin{itemize} \\item New observations of \\J conducted with the GMRT and VLA confirm its morphology consisting of only two narrow lobes; no trace of any outer cocoon is therefore detected. We model the newly accessed radio data using the \\D algorithm of \\citet{mach07} to determine --- for a given set of observables --- the main physical parameters of the jets inflating the lobes and of their environment. Different model fits performed for the observed radio structure imply relatively young age of the source $\\sim 35$\\,Myr, its relatively high expansion velocity $\\sim 0.2\\,c$, large kinetic power $\\sim 4 \\times 10^{45}$\\,erg\\,s$^{-1}$, and confirm particularly low-density environment $\\sim 2 \\times 10^{-31}$\\,g\\,cm$^{-3}$. The kinematic estimates of the age and of the advance velocity are fully consistent with the values resulting from the \\D modeling. \\item We find that it is possible to chose a realistic set of the model parameters for which the hypothetical outer lobes of \\J are old enough so that their expected radio surface brightness is substantially below the rms noise level of the available radio maps. We speculate that future high-sensitivity VLA and Low Frequency Array observations, or X-ray studies using facilities such as XMM-{\\it Newton} or {\\it Suzaku}, may be more promising for investigating a presence of the hypothetical outer cocoon in \\J. \\item The extremely low density of the gas surrounding the observed lobes in \\J is consistent with the mean density of the baryonic matter in the Universe. This suggests that if the recurrent jet activity in \\J is not the case, the source has to be located instead in a real void of the galaxy and matter distribution. This conclusion is partly supported by the galaxy counts performed in the vicinity of the studied object. If correct, it further implies that the giant lobes in \\J are expected anyway to modify dramatically their surroundings by driving strong shocks (Mach numbers $>20$) and heating the gas located at the outskirts of the filamentary galactic distribution (`warm-hot phase of the intergalactic medium'). \\item Finally, we also find that the energetic requirements for the source (with a given mass of the central supermassive black hole estimated in this paper) are severe in terms of the total jet power and the total energy deposited by the outflows far away from the central engine. Rough estimates presented imply therefore that this central engine --- with little manifestation of the ongoing accretion-related activity at present --- was once a luminous AGN, radiating at the maximum (Eddington) rate. This provides interesting constraints on the AGN duty cycle in \\emph{isolated} early-type galaxies. \\end{itemize}"
    },
    "1107/1107.1483_arXiv.txt": {
        "abstract": "Disk galaxies at high redshift have been predicted to maintain high gas surface densities due to continuous feeding by intense cold streams leading to violent gravitational instability, transient features and giant clumps. Gravitational torques between the perturbations drive angular momentum out and mass in, and the inflow provides the energy for keeping strong turbulence. We use analytic estimates of the inflow for a self-regulated unstable disk at a Toomre stability parameter $Q\\sim 1$, and isolated galaxy simulations capable of resolving the nuclear inflow down to the central parsec. We predict an average inflow rate $\\sim\\!10\\,\\sy$ through the disk of a $10^{11}\\,\\msun$ galaxy, with conditions representative of $z\\!\\sim\\!2$ stream-fed disks. The inflow rate scales with disk mass and $(1+z)^{3/2}$. It includes clump migration and inflow of the smoother component, valid even if clumps disrupt. This inflow grows the bulge, while only a fraction $\\gsim\\!10^{-3}$ of it needs to accrete onto a central black hole (BH), in order to obey the observed BH-bulge relation. A galaxy of $10^{11}\\msun$ at $z\\!\\sim\\!2$ is expected to host a BH of $\\sim\\! 10^8\\msun$, accreting on average with moderate sub-Eddington luminosity $L_\\mathrm{X}\\!\\sim\\!10^{42-43}\\,\\ergs$, accompanied by brighter episodes when dense clumps coalesce. We note that in rare massive galaxies at $z\\!\\sim\\!6$, the same process may feed $\\sim\\!10^9\\,\\msun$ BH at the Eddington rate. High central gas column densities can severely obscure AGN in high-redshift disks, possibly hindering their detection in deep X-ray surveys. ",
        "introduction": "The growth of massive black holes (BH) and the associated Active Galactic Nuclei (AGN) are commonly assumed to be driven by gas inflows from galaxy mergers \\citep{DM05,H06}. However, baryons are fed into high-redshift galaxies through cold streams, in which the contribution of smooth gas and small clumps is larger than that of mergers of mass ratio greater than 1:10 \\citep{dekel09, brooks09}.   Cold accretion and high gas fractions in high-redshift disks trigger a violent instability, often characterized by giant star-forming clumps. Using analytic estimates (Section~2) and simulations (Section~3), we address the possibility that the gas supply for BH growth is provided by instability-driven inflows, and evaluate the obscuration of AGN in this model (Section~4).    \\newpage ",
        "conclusions": "Violent gravitational instability in high-redshift disk galaxies naturally leads to the growth of a bulge and a central BH; it can both feed an AGN and obscure it. This stems from the developing picture where many star-forming galaxies at high redshift are rotating disks with high gas fractions continuously fed by cold streams. Such disks develop instability that involves transient perturbations and giant clumps. Gravitational torquing among these perturbations lead to mass inflow, which provides energy for maintaining the strong turbulence required for self-regulated instability at $Q\\! \\sim \\! 1$. The cosmological inflow rate along cold streams sets an upper limit, but the actual inflow rate toward the nucleus is determined by the disk instability on galactic scales.   High-redshift disks are very different from nearby spirals, the key difference being the much higher gas fraction. The instability then operates on short dynamical timescales, and drives an intense inflow through the disk at the level of $\\sim \\!10\\sy M_{\\rm b,11} (1+z)_3^{3/2}$. In low-redshift disks, the instability is a secular process limited to non-axisymmetric modes with considerably weaker torquing: bars are invoked in some AGN models \\citep{fanidakis, begelman} but need about ten rotation periods to convey some gas inward \\citep{BCS05}. Gravitational torquing in high-$z$ disks is much more efficient and involves richer gas reservoirs.  \\medskip  Disk instability at $z \\! \\sim \\! 2$ can thus funnel half of the disk gas toward the center in 2~Gyr. This is similar to the mass inflow in a major merger \\citep{H06}, but spread over a ten times longer period, resulting in a lower average AGN luminosity, with higher duty cycle, and high obscuration. This process could dominate, as the gas mass in the relatively smooth cosmological streams is larger than that associated with mergers -- less than 10\\% of the cosmological infall is in major mergers \\citep{dekel09, lhuillier11}. The main prediction is thus that many high-$z$ AGNs should be hosted by star-forming disk galaxies, composed of clumpy disks and growing spheroids. Many would be moderately luminous, detectable by {\\it Chandra} deep surveys at $z\\!=\\!2$ if un-obscured, but the high expected obscuration makes their detectability dubious, especially for galaxies below $10^{11}\\,\\msun$. These BHs can grow inside classical bulges \\citep{EBE08b}.   Observationally, the fraction of X-ray-detectable AGN indeed declines with decreasing stellar mass \\citep{xue10}. Many AGN are highly absorbed, individually undetectable in X-rays \\citep[e.g.,][]{daddi07}, and many of these obscured AGN have moderate intrinsic luminosities and lie in star-forming galaxies \\citep{J11}. Host galaxies of moderate-luminosity AGNs do not show an excess of major-merger signatures \\citep{grogin05,gabor09} and often have disky morphologies \\citep{schawinski11}. Small mergers may also help feeding the BH while leaving the disk intact \\citep{shankar10}, but they are a component of the cosmological streams, and can be considered to be part of the inflow within the disk.  \\medskip  The violent instability phase should end when the cosmological accretion rate declines and the system becomes stellar dominated. We expect less massive galaxies to remain unstable for longer times, because they retain higher gas fractions --- one of the aspects of the downsizing of star formation \\citep[e.g.,][]{J05}. This can result from the regulation of gas consumption in smaller galaxies \\citep{DekelSilk86, KD11} and from the continuation of cold accretion to lower redshift for lower-mass halos \\citep{DB06}. It induces an inverse gradient of gas fraction with galaxy mass (as observed, \\citealt{kannappan04}) and a downsizing in gravitational instability which could result in downsizing of BH growth. This is a longer growth phase into later redshifts in lower-mass galaxies, with $\\sim \\! 10^{11}\\msun$ galaxies growing BHs mostly at $z > 1$ in their unstable disk phase, while clumpy disks of $\\sim \\! 10^{10}\\msun$ may show AGN activity even after $z \\! \\sim \\! 1$."
    },
    "1107/1107.3165_arXiv.txt": {
        "abstract": "We explore the potential cumulative energy production of stellar mass black holes in early galaxies. Stellar mass black holes may accrete substantially from the higher density interstellar media of primordial galaxies, and their energy release would be distributed more uniformly over the galaxy, perhaps providing a different mode of energy feedback into young galaxies than central supermassive black holes. We construct a model for the production and growth of stellar mass black holes over the first few gigayears of a young galaxy. With the simplifying assumption of a constant density of the ISM, $n \\sim 10^4 - 10^5$ cm$^{-3}$, we estimate the number of accreting stellar mass black holes to be $\\sim 10^6$ and the potential energy production to be as high as $10^{61}$ergs over several billion years. For densities less than $10^5$ cm$^{-3}$, stellar mass black holes are unlikely to reach their Eddington limit luminosities. The framework we present could be incorporated in numerical simulations to compute the feedback from stellar-mass black holes with inhomogeneous, evolving interstellar media.  Keywords: Galaxies: evolution - Galaxies: stellar content - Stars: massive - Black hole physics - Accretion, accretion disks - Radiative transfer ",
        "introduction": "Most research on the early formation and growth of black holes has focused on the important issues of the growth and feedback of intermediate and supermassive black holes; relatively little attention has been paid to the effects of stellar-mass black holes and their potential energy production. Estimates of the number of stellar mass black holes presently in the Galaxy are as high as $10^{8}$ to $10^{9}$ \\citep{AK02,CDK03}, suggesting that the nearest is only parsecs away, but accreting feebly from the modern interstellar medium.  The cumulative energy production of $10^{8}$ $10M_{\\odot}$ black holes could rival that of a single central $10^{9}M_{\\odot}$ black hole.  The first stars that formed were extremely low metallicity and arguably high mass \\citep{BCL02} Population III stars. These stars may have only produced intermediate-mass or supermassive black holes, if any, but turbulence may broaden the initial mass function so that lower mass stars and black holes form even in the first generation \\citep{Clark11,Prieto11}. The next generation of stars, Population II, were slightly higher in metallicity and would form stellar mass black holes. This epoch might have begun at redshift in excess of 15 \\citep{BH06} during the era of reionization, when the interstellar medium number density, $n_{ism}$, might still be large. In the early formation of these galaxies, their interstellar medium (ISM) density was probably much higher than it is currently, perhaps $n_{ism} \\sim 10^{4}$cm$^{-3}$ \\citep{Greif08}. With ISM densities this high, stellar-mass black holes could potentially accrete substantial amounts of matter over the course of the early evolution of galaxies and emit energy back into the galaxy and intergalactic medium. Since the number of stellar mass black holes could be as high as $10^{8}$, their cumulative energy production could be extremely high, given a favorable accretion and feedback rate.  Since stellar mass black holes would be distributed more evenly over the galaxy than a single supermassive black hole, the energy feedback may impact the evolution of a galaxy in different ways than would a singular source. Stellar-mass black holes also deliver feedback in a different mode than supernovae, the energy of which is primarily kinetic, and, unlike supernovae, black holes persist. The rate of star formation in galaxies is about a factor of 100 less than provided by basic estimates \\citep{K98,S08}. This suggests a feedback process that has not yet been identified. The process we outline here may contribute to this regulation of star formation in early galaxy evolution.  We apply the prescriptions of black hole accretion with feedback from \\citep{MBCO09} and \\citet{PR11} to the stellar mass black hole case to approximate the energy production in early galaxies. With these estimations, we find that the potential energy production of the stellar mass case could be on the order of $10^{60}$ ergs over $\\sim 2$ Gyr. Section 2 gives an outline of our model for the birth and growth by accretion of stellar mass black holes. Section 3 addresses the rate of accretion of stellar mass black holes including the effects of radiative feedback and gives estimates of the number of black holes, their luminosity, and their cumulative energy liberated as a function of time. Section 4 discusses the possible constraints of the Eddington limit and \\S5 gives our conclusions. ",
        "conclusions": "A plethora of stellar-mass black holes accreting in young galaxies could be significant source of radiant energy. Unlike the feedback from stars and supernovae that are one-time, rather short term events in the history of a galaxy, individual stellar-mass black holes accumulate and continue to radiate as long as the conditions promote accretion. A principal factor is the ambient density of the ISM that must remain sufficiently high.  We have shown that for predicted ambient densities in the range $10^4 - 10^5$ cm$^{-3}$, the condition for appreciable accretion is met in the context of Bondi/Hoyle accretion limited by radiative feedback. In this case, there is potentially a significant amount of accretion-radiated energy available.  We derive the total energy production of stellar mass black holes accreting from a high-density primordial ISM using the classical Bondi spherically symmetric accretion rate scaled with the value for the mean accretion rate affected by feedback given in \\citet{PR11}. Since stellar mass black holes will be distributed in proportion to star formation, their energy production will be emitted more uniformly over the galaxy than a single supermassive black hole. This may have implications for early galaxy evolution as large scale galaxies start to form during the era of reionization. We also show that given the estimated values used in our calculations, the stellar-mass black holes are not likely to reach the Eddington limit during the time span we have used here, a few Gyr.  For our analytic calculations, the ISM density, $n_{ism}$, has been taken to be constant in space and over the time spans of interest. In reality, the density will vary througout the galaxy and as the galaxy evolves over time, the density will decrease. The motion of the black holes with respect to the background gas could also be an important limiting factor as the motion approaches or exceeds supersonic. This motion could affect the nature of the feedback process itself, but, to the best of our knowledge, this has not been explored in any detail.  Due to these effects, our calculations may be more optimistic than a more realistic case that accounts for density fluctuations and evolution and for black hole motion. Some account of these factors might be taken by replacing the density of the ISM in our formulation, $n_{ism}$ with $f_f n_{ism}$ where $f_f$ is a filling factor corresponding to a given density. The ambient density may be affected by whether star formation is dominated by mergers or by cold inflow at the epochs of interest \\citep{S08}. We have, however, given a general framework by which the accumulating number of stellar-mass black holes, their luminosity at a given epoch, and their accumulated feedback energy could be incorporated in a numerical simulation that followed the fluctuations and density evolution more realistically. A simulation that included the distributed feedback effects of accumulating, accreting stellar-mass black holes might give significantly different results than one that only included stars, supernovae and one or a few intermediate-mass or supermassive black holes.  As we were finishing work on this paper, the paper by \\citet{Mirabel11} appeared that considers the possible role of stellar-mass black holes in high-mass X-ray binaries in young galaxies. This scenario has the advantage that the accretion is driven by mass transfer from a companion star and does not depend on the vagaries of the ISM. On the other hand, this proposition does depend on the binary fraction of massive stars in the early Universe, a rather uncertain quantity. \\citet{Mirabel11} seem to assume that all stellar-mass black holes formed in these early epochs are in binary systems. We note that these high-mass X-ray binaries live for a short time \\citep[perhaps 0.02 Gyr,][]{Mirabel11} and hence each binary system is a delta-function contribution on the time scales we consider here. When the secondary dies, the primary black hole (or a secondary one if it forms) would then be subject to the sort of radiative-feedback limited Bondi accretion we consider here. Both single and binary black holes should be considered in complete models of the evolution of young galaxies.  While we have concentrated on epochs $\\sim$ Gyr, the existence and impact of stellar-mass black hole accretion and feedback may be pertinent to the earliest phases of star and galaxy formation if, for instance, turbulence allows the formation of a broader initial mass function and hence lower mass stars already in the first dark matter halos where star formation is thought to have first begun \\citep{Clark11,Prieto11}. Whereas a single very massive star explosion by pair-instability releasing $\\sim 10^{53}$ ergs is likely to unbind a mini-halo, the explosion of less massive ``normal\" supernovae and the accretion and feedback of stellar-mass black holes may yield a different behavior than many current simulations envisage. Supernovae may blow fountains rather than totally disrupting the baryonic content of the halos and feedback from accreting black holes may alter state of ionization of the gas. Both types of events may drive further turbulence. If stellar-mass black holes are present from the first epochs of star formation, then their role must be considered in the subsequent mergers that build more massive galaxies. Aside from the effects of accretion and feedback, any stellar-mass black holes that form and remain in mini halos may themselves merge through dynamical friction and help to promote the growth of larger-mass black holes during the era of rapid galaxy merging. The tendency to move supersonically with respect to the ambient gas that will tend to limit the accretion may promote this dynamical friction and merging.  An important aspect that we have not yet explored in depth is the observational consequence of our model. This might be addressed by a simulation incorporating stellar-mass black hole feedback. A related issue might be to discriminate the signal of a host of stellar-mass black holes from single, larger-mass black holes. One means to do this might be to make both radio and X-ray observations of young galaxies. There appears to be a nearly universal relation between the radio and X-ray luminosity of black holes such that $L_{radio} \\propto L_X^{0.7}$, where both luminosities are presumed to scale monotonically with black hole mass. If this holds for the conditions we address here, then an ensemble of stellar-mass black holes should have a larger radio flux for a given X-ray flux than a single black hole of the same total mass. In this context, \\citet{Jia11} have examined the evidence for the growth of black holes in local starburst galaxies. They find an anomalous increase in X-ray flux that they argue could be associated with the growth of a black hole of $10^4$\\m\\ by accretion. We suggest there might be an ensemble of $10^3$ smaller-mass black holes that might provide the same X-ray flus, in which case the radio luminosity would be relatively large.  Combining Eqns \\ref{mdot2} and \\ref{Fnumber} we estimate the accretion rate for fiducial parameters to be $\\sim 10^{-9}$ \\m\\ yr$^{-1}$. This accretion rate corresponds roughly to that needed to induce the accretion disk limit-cycle instability \\citep{CGW82,LFP85} and thus to produce a black hole X-ray nova analogous to AO620-00 and related events \\citep{TL95}. The X-ray flux from these systems might thus come in flares near the Eddington limit lasting for months with quiescent periods of years to decades. In an active star-forming galaxy at a red shift of about 2 with $\\sim 10^6$ stellar mass black holes, $\\sim10^4$ of them might be in outburst at any given time. Any such bursts may be redshifted into bands that are heavily extincted and hence difficult to observe directly, but this aspect is worth considering more carefully.  \\citet{K11} have argued that black holes do not correlate with disks and that they correlate little, if at all, with pseudobulges. It would be interesting to consider whether or not stellar-mass black holes affect the early evolution of disk-grown pseudobulges.  There is currently a problem understanding the rapid early growth of dust in young galaxies. Perhaps the ensemble of stellar-mass black holes contributes to dust formation by, for instance, providing numerous and prolonged concentrations of ISM density at the boundary of radiative feedback regions.  Finally we note that our estimates here suggest that a contemporary stellar mass black hole that found itself in a dense clump of molecular material, $n \\sim 10^5$ cm$^{-3}$, might be observably luminous $\\sim 10^{37}$ erg s$^{-1}$ for fiducial parameters at this density. Dense clumps like this are estimated to fill a volume in the Galaxy of about $4\\times10^4$ pc$^3$ (N. J. Evans, private communication). With a volume within 100 pc of the Galactic plane of $\\sim 10^{11}$ pc$^3$, this represents a filling factor for dense molecular clumps of $\\sim 4\\times10^{-7}$. With estimates of $10^6 - 10^8$ black holes in the Galaxy, the estimated number of black holes within such dense clumps ranges from negligible to a few."
    },
    "1107/1107.3812_arXiv.txt": {
        "abstract": "Higher order Laguerre--Gauss (LG) beams have been proposed for use in future gravitational wave detectors, such as upgrades to the Advanced LIGO detectors  and the Einstein Telescope, for their potential to reduce the effects of the thermal noise of the test masses.  This paper details the theoretical analysis and simulation work carried out to investigate the behaviour of LG beams in realistic optical setups, in particular the coupling between different LG modes in a linear cavity. We present a new analytical approximation to compute the coupling between modes, using Zernike polynomials to describe mirror surface distortions. We apply this method in a study of the behaviour of the LG$_{33}$ mode within realistic arm cavities, using measured mirror surface maps from the Advanced LIGO project. We show mode distortions that can be expected to arise due to the degeneracy of higher order spatial modes within such cavities and relate this to the theoretical analysis. Finally we identify the mirror distortions which cause significant coupling from the LG$_{33}$ mode into other order 9 modes and derive requirements for the mirror surfaces. ",
        "introduction": "The sensitivities of second generation gravitational wave detectors such as Advanced LIGO and Advanced Virgo are expected to be limited by the thermal noise of the test masses within a significant range of signal frequencies around 100 Hz~\\cite{Rowan05}. To reach even better sensitivities, it has been proposed to use laser beams with an intensity pattern other than that of the fundamental Gaussian beam to reduce the effects of this thermal noise~\\cite{Vinet07, Mours06}.  A beam whose intensity is distributed more homogeneously over the mirror surface, for the same clipping losses, benefits from a more effective averaging over the mirror surface distortions caused by thermal effects~\\cite{Vinet09}. The specific advantage of using higher order LG modes, as opposed to mesa~\\cite{Ambrosio03} and conical~\\cite{Bondarescu08} beams, is that they are compatible with spherical mirrors as currently used in GW detectors and other high precision optical setups. Research into the potential of the LG$_{33}$ mode in gravitational wave detectors has been carried out using numerical simulations and table-top experiments~\\cite{Chelkowski09, Fulda10}. The sensing and control signals for an  LG$_{33}$ beam were found to perform as well as for the fundamental mode in all aspects examined and the LG$_{33}$ behaved as expected in short linear and triangular optical cavities. However, an optical cavity resonant for a higher--order Gaussian mode is degenerate so that a number of modes can resonate at the same time. This is a fundamental difference to a well designed cavity for the fundamental Gaussian mode, in which any resonant enhancement of other modes can be suppressed. This degeneracy can potentially cause additional optical losses. Simulations have shown that the use of the LG$_{33}$ beam, compared to the fundamental mode, LG$_{00}$, could result in a significant contrast defect at the dark fringe~\\cite{Miller,Galimberti,Yamamoto}. It is the aim of this paper to investigate how mirror surface distortions affect the purity of an LG$_{33}$ beam in high finesse cavities, by analytical calculation and numerical simulation.  We will focus on the direct coupling from a distorted mirror and how this affects the mode content in a linear cavity, with our final aim to produce specifications for the mirror surfaces. ",
        "conclusions": "We have investigated the coupling which occurs when Laguerre-Gauss modes are incident on a mirror with surface distortions.  Taking an analytical approach we used Zernike polynomials to represent mirror surface distortions and derived an approximate equation for the significant coupling when an LG mode is reflected from a surface defined by a particular Zernike polynomial.  This derivation resulted in a condition for significant coupling, $m=|l-l'|$, where $m$ is the azimuthal index of the Zernike polynomial and $l$ and $l'$ are the azimuthal indices of the incident and coupled modes respectively.  This is a significant result as it allows us to predict which order 9 modes will be largely coupled by particular Zernike polynomials and suggest which modes will have large amplitudes in the arm cavities  We investigated the performance of LG$_{33}$ in high finesse cavities by simulation with Advanced LIGO mirror maps.  This illustrated the degraded purity of the circulating beam in realistic cavities due to higher order mode degeneracy.  The results were then analysed by looking at the Zernike polynomials representing our example mirror map.  The analysis and results were consistent with the predictions made from Eq.~\\ref{eq:k_eq}.  This suggested that astigmatism was causing a significant amount of coupling, particularly into the LG$_{41}$ and LG$_{25}$ modes.  This was confirmed when the cavity was simulated again with the astigmatism removed from the mirror map and we observed a dramatic increase in the LG$_{33}$ mode purity.  The analytical description enabled us to identify the specific Zernike polynomials which cause large couplings as well as the LG modes which would dominate as a result.  Using this we were able to derive certain requirements for our example mirror map, ETM08, in terms of limits on the amplitudes of the Zernike polynomials Z$_{2}^{2}$, Z$_{4}^{2}$ and Z$_{4}^{4}$ (table~\\ref{table:limits}).  Using this map the resulting circulating beam impurity was found to be 815\\,ppm, a significant reduction from the original impurity of 0.114.  This investigation has demonstrated that a high beam purity is achievable using an LG$_{33}$ beam when modifications are made to the low order Zernike polynomials in Advanced LIGO mirrors.  Implementing the LG$_{33}$ beam in gravitational wave detectors will be challenging as we require very small amplitudes on these lower order polynomials.  We should also consider that the example mirror surfaces considered here refer to uncoated substrates.  The coating process is likely to add to the lower order features in the mirror surfaces.  However, using this analytical approach we can derive specific requirements for the mirror surfaces leading to designs for suitable mirrors for these higher order beams."
    },
    "1107/1107.2032.txt": {
        "abstract": "{We report the detection of CoRoT-18b, a massive hot jupiter transiting in front of its host star with a period of $1.9000693 \\pm 0.0000028$~days. This planet was discovered thanks to photometric data secured with the CoRoT satellite combined with spectroscopic and photometric ground-based follow-up observations. The planet has a mass \\mp$\\,  = 3.47 \\pm 0.38$\\,\\Mjup, a radius $R_{\\rm p} = 1.31\\pm0.18$\\,\\Rjup, and a density $\\rho_{\\rm p} = 2.2 \\pm 0.8$\\,g/cm$^3$. It~orbits a G9V star with a mass $M_\\star = 0.95\\pm0.15$\\,M$_\\odot$, a radius $R_\\star = 1.00\\pm0.13$\\,R$_\\odot$, and a rotation period $P_{\\rm rot} = 5.4\\pm0.4$\\,days. The age of the system remains uncertain, with stellar evolution models pointing either to a few tens Ma or several Ga, while gyrochronology and lithium abundance point towards ages of a few hundred Ma. This mismatch potentially points to a problem in our understanding of the evolution of young stars, with possibly significant implications for stellar physics and the interpretation of inferred sizes of exoplanets around young stars. We detected the Rossiter-McLaughlin anomaly in the \\cible\\ system thanks to the spectroscopic observation of a transit. We measured the obliquity $\\psi = 20^{\\circ} \\pm 20^{\\circ}$ (sky-projected value $\\lambda = -10^{\\circ} \\pm 20^{\\circ}$), indicating that the planet orbits in the same way as the star is rotating and that this prograde orbit is nearly aligned with the stellar~equator. ",
        "introduction": "Out of the $\\sim550$ exoplanets known to date, more than 100 transit their parent stars as seen from the Earth. This particular configuration allows numerous key studies, including accurate radius, mass, and thus density measurements (see, e.g., Winn~\\cite{winn10a} for a review), atmospheric studies in absorption through transits and in emission through occultations (e.g. Vidal-Madjar et al.~\\cite{avm03};  Wheatley et al.~\\cite{wheatley10}), dynamic analyses through possible timing variations (e.g. Holman et al.~\\cite{holman10}), or spin-orbit alignment measurements thanks to the  Rossiter-McLaughlin effect (e.g. Bouchy et al.~\\cite{bouchy08}). The power of these analyses incited numerous search surveys for transiting planets. Most of them were discovered in the last five years, and the detection rate is still increasing.  The CoRoT space mission (\\textit{COnvection ROtation and planetary Transits}, Baglin et al.~\\cite{baglin09}) was launched on 2006 December 27. Based on a 27-cm telescope and a $2.8\\degr \\times 2.8\\degr$-field camera, it is designed to study asteroseismology and detect transiting exoplanets. The satellite allows several thousand stars ($V = 12 - 16$) to be continuously observed for up to 150~days with a high photometric accuracy. CoRoT is thus well adapted to detecting transiting planets with small radii, such as CoRoT-7b (L\\'eger et al.~\\cite{leger09}; Queloz et al.~\\cite{queloz09}), or on long orbital periods, such as CoRoT-9b (Deeg et al.~\\cite{deeg10}). It can also detect hot jupiters, such as \\cibleb. We report its discovery~here.  We describe in Sect.~\\ref{sect_corot_obs} the CoRoT observations and the transit detection of the planetary candidate. Then, we present in Sect.~\\ref{sect_ground_obs} the ground-based follow-up observations that were needed to establish the planetary nature of the event detected by CoRoT and also to characterize this planetary system. The analysis of the whole dataset and the results are presented in Sect.~\\ref{sect_analysis}, before conclusion in Sect.~\\ref{sect_concle}. ",
        "conclusions": "\\label{sect_concle}  We reported the detection of the 18$^{\\rm th}$ transiting exoplanet detected by the CoRoT project. This giant planet was discovered thanks to the high-accuracy, continuous photometry obtained by the CoRoT satellite and the photometric and spectroscopic follow-up performed on ground-based telescopes. \\cibleb\\ is a massive hot jupiter orbiting a faint G9V star. Its mass is \\mp\\,$  = 3.47 \\pm 0.38$\\,\\Mjup, and its radius $R_{\\rm p} = 1.31\\pm0.18$\\,\\Rjup, implying a density $\\rho_{\\rm p} = 2.2 \\pm 0.8$\\,g/cm$^3$. The period of the circular orbit is $1.9000693 \\pm 0.0000028$~days. It is known with an accuracy better than 0.25~seconds and the mean mid-transit epoch with an accuracy of 20~seconds. The mass of the host star  is $M_\\star = 0.95\\pm0.15$\\,M$_\\odot$ and its radius $R_\\star = 1.00 \\pm0.13$\\,R$_\\odot$. The parameters of this system are summarized in Table~\\ref{starplanet_param_table}.  The parameters of \\cibleb\\ are similar to those of CoRoT-2b (Alonso~\\cite{alonso08}; Gillon et al.~\\cite{gillon10}), and to a lesser extent to those of CoRoT-11b (Gandolfi~\\cite{gandolfi10}) and CoRoT-17b (Csizmadia et al.~\\cite{csizmadia11}), the other massive hot jupiters found with CoRoT. Interestingly \\cibleb\\ is found to be either particularly young (a few tens of Ma) or old ($>4$\\,Ga) from stellar evolution models matching the star's effective temperature and inferred density, but according both to its lithium abundance and to its relatively fast rotation, it would be expected to be modestly young (several hundred Ma). This mismatch potentially points to a problem in our understanding of the evolution of young stars, with possibly significant implications for stellar physics and the interpretation of inferred sizes of exoplanets around young stars. In addition, \\cibleb\\ and CoRoT-2b are the only known planets (transiting or not) orbiting a fast-rotating, G-type or cooler star.  The orbit of \\cibleb\\ is prograde, with a spin-orbit angle $\\psi = 20^{\\circ} \\pm 20^{\\circ}$ (sky-projected value $\\lambda = -10^{\\circ} \\pm 20^{\\circ}$), hence an obliquity in agreement with 0. Schlaufman~(\\cite{schlaufman10}) and Winn et al.~(\\cite{winn10b}) have shown that misaligned planets tend to orbit hot stars. With an effective temperature of $5440 \\pm 100$\\,K  for the host star, the \\cible\\ system supports this trend. Exceptions to this trend are increasing, however (Moutou et al.~\\cite{moutou11}; Brown et al.~\\cite{brown11}).  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.525]{hebrardf13.pdf} \\caption{Histogram of the number of known transiting planets as a function of their mass (data from http://exoplanet.eu). } \\label{fig_histo} \\end{center} \\end{figure}   H\\'ebrard et al.~(\\cite{hebrard10}; \\cite{hebrard11}) have hypothesized that most of the massive planets are prograde and moderately but significantly misaligned, whereas the less massive planets are distributed in two thirds of the prograde, aligned systems and one third of the strongly misaligned systems. Being prograde and nearly aligned, \\cibleb\\ is at the upper limit of the low-mass range. Interestingly, a limit is also apparent in the mass distribution of known transiting planets. This is shown in Fig.~\\ref{fig_histo}, where one can see a decreasing abundance of planets with increasing planetary mass up to 4.5\\,\\Mjup. No transiting planets are known in the range \\mp$  = [4.5-7]$\\,\\Mjup, and a few are known with \\mp$>7$\\,\\Mjup. It is difficult to imagine that a bias is provoking the lack of massive planets, as they are easier to detect. This different distribution suggests that planets below  4.5\\,\\Mjup\\ could have a different nature or history than those above 7\\,\\Mjup. This is reinforced by the different obliquity~distribution."
    },
    "1107/1107.3338.txt": {
        "abstract": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We report evidence of ordered orbital motion in luminous star-forming galaxies at $z\\sim1.3$. We present integral field spectroscopy (IFS) observations, performed with the OH Suppressing InfraRed Imaging Spectrograph (OSIRIS) system, assisted by laser guide star adaptive optics on the Keck telescope, of 13 star-forming galaxies selected from the WiggleZ Dark Energy Survey. Selected via ultraviolet and \\OII~emission, the large volume of the WiggleZ survey allows the selection of sources which have comparable intrinsic luminosity and stellar mass to IFS samples at $z>2$.  Multiple 1--2 kpc size sub-components of emission, or `clumps', are detected within the \\halpha spatial emission which extends over 6--10 kpc in 4 galaxies, resolved compact emission ($r<3$ kpc) is detected in 5 galaxies, and extended regions of \\halpha emission are observed in the remaining 4 galaxies. We discuss these data in the context of different snapshots in a merger sequence and/or the evolutionary stages of coalescence of star-forming regions in an unstable disk. We find evidence of ordered orbital motion in galaxies as expected from disk models and the highest values of velocity dispersion ($\\sigma>100$ km s$^{-1}$) in the most compact sources. This unique data set reveals that the most luminous star-forming galaxies at $z>1$ are gaseous unstable disks indicating that a different mode of star formation could be feeding gas to galaxies at $z>1$, and lending support to theories of cold dense gas flows from the intergalactic medium.\\\\ ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% An open question remains regarding the accretion of gas and build up of stars in massive galaxies at the peak epoch of cosmic star formation. The traditional theories of isolated disk-formation resulting from gas accretion from the halo \\citep{1962ApJ...136..748E} have been challenged by recent observations and theories at high redshift. In new models of disk formation at high redshift, cooling flows supply gas directly to the centres of galaxies \\citep{2005MNRAS.363....2K,Dekel:2009ys}. Direct observations of the gas motions in galaxies are crucial to test these new models and ultimately to disentangle how galaxies form at high redshift.  Morphological studies made possible by Hubble Advanced Camera for Surveys (ACS), Wide Field and Planetary Camera (WFPC) and other deep imagers have provided the first clues to the processes governing high-redshift galaxy formation. High-redshift galaxies exhibit different morphologies to local galaxies, with higher fractions of both compact and irregular galaxies \\citep{Glazebrook:1995vn, 1996MNRAS.279L..47A}. This deviation from the Hubble Sequence indicates that different mechanisms may be driving star formation and galaxy evolution at earlier cosmic times. Initial interpretations of the complex structures seen in the Hubble Ultra Deep Field (HUDF) favor higher rates of major mergers, a natural consequence of hierarchical merging models in a $\\Lambda$CDM Universe (e.g., \\citealt{2003AJ....126.1183C, 2006ApJ...636..592L,  2008ApJ...677...37O,  Engel:2010fk}). Yet, despite large dense regions in these galaxies, observations of their stellar populations reveal galactic properties reminiscent of local disk galaxies \\citep{2005ApJ...627..632E}. In this case the large dense regions or `clumps' of star formation may be related to HII regions found in spiral galaxies but of much larger size \\citep{Jones:2010uf,2010Natur.464..733S}. Models have been developed to explain the clumpy structures in disks and their evolution into galactic structures seen at $z=0$ (e.g., \\citealt{1999ApJ...514...77N, 2004A&A...413..547I, 2007ApJ...670..237B, Elmegreen:2008fk}). In these models, turbulent gas collapses from Jeans instabilities resulting in massive star-clusters that, on a timescale of 1 Gyr, coalesce to the centre of the galaxy to form a pseudo-bulge.  %subsection{IFU RESULTS} Morphological studies alone do not provide enough information to distinguish between the two kinematically distinct formation theories of merging and isolated disk formation at these redshifts. Given the complex structure of the galaxies it is difficult to derive the major axis for slit positioning making spectroscopic observations less effective in discerning rotation signatures (e.g. \\citealt{2004ApJ...612..122E,2006ApJ...646..107E}). Spatially-resolved spectroscopy is therefore essential for disentangling kinematics and star-formation histories of galaxies at high redshift.  The past five years have seen a great number of integral field spectroscopy (IFS) studies of the nebular emission line, \\halphans, in $z>2$ star-forming galaxies (e.g., \\citealt{2006Natur.442..786G, 2008ApJ...687...59G,2009ApJ...706.1364F, 2009ApJ...697.2057L, 2007ApJ...658...78W,2009ApJ...699..421W, 2009A&A...504..789E, 2008A&A...479...67N}). These redshifts are favored due to the large number of Lyman break galaxies (LBGs) and sub-millimeter galaxies (SMGs) being discovered by methods tailored for high redshift as well as the accessibility of \\halpha in the near-infrared (NIR).  \\begin{table*}%\\footnotesize \\begin{minipage}{\\textwidth} \\caption{OSIRIS observations of WiggleZ galaxies} \\begin{tabular*}{\\textwidth}{@{\\extracolsep{\\fill}}llrrcccccc} \\hline {WiggleZ ID} & {ID} & {R.A.\t} & {Dec. \t} & {z$^{a}$\t} & {Obs.} & {$t_{\\mathrm{exp.}}$    } & {Exp.    } & {Filter\t} & {Scale\t} \\\\ {} & {} & {(J2000.0)\t} & {(J2000.0)\t} & {} & {Date} & {(s)} & {Seq. } & {} & {(mas)} \\\\ \\vspace{0.25mm} \\\\ \\hline \\multicolumn{10}{c}{Detections}\\\\ \\hline \\vspace{0.25mm} \\\\ R03J032450240$-$13550943\t& WK0912\\_13R &  03:24:50.240  & $-$13:55:09.429  &  1.288\t&  2009/12\t& 3600 & 4 $\\times$ 900 &  Hn1\t& 50  \\\\ S15J145355248$-$00320351\t& WK1002\\_61S &  14:53:55.247  & $-$00:32:03.501  &  1.305  &  2010/02  & 5400 & 6 $\\times$ 900 &  Hn1  & 50 \\\\ R01J005822757$-$03034040\t& WK0909\\_02R &  00:58:22.757  & $-$03:03:40.400  &  1.362\t&  2009/09\t& 5400 & 9 $\\times$ 600 &  Hn2\t& 50  \\\\ R03J032206214$-$15443471\t& WK0909\\_06R &  03:22:06.214  & $-$15:44:34.711  &  1.461\t&  2009/09\t& 8100 & 9 $\\times$ 900 &  Hn3\t& 50  \\\\  S15J144102444+05480354\t& WK0905\\_22S &  14:41:02.443  &  05:48:03.540  &  1.282\t&  2009/05\t& 3600 & 4 $\\times$ 900 &  Hn1\t& 50  \\\\ R00J232217805$-$05473356\t& WK0912\\_01R &  23:22:17.805  & $-$05:47:33.560  &  1.297\t&  2009/12\t& 3600 & 4 $\\times$ 900 &  Hn1\t& 50  \\\\ S09J091517481+00033557\t& WK1002\\_41S &  09:15:17.483  &  00:03:35.570  &  1.334  &  2010/02  & 7200 & 8 $\\times$ 900 &  Hn1  & 50  \\\\ S00J233338383$-$01040629\t& WK0809\\_02S &  23:33:38.386  & $-$01:04:06.290  &  1.452\t&  2008/09\t& 5400 & 6 $\\times$ 900 &  Hn3\t& 50  \\\\  S11J101757445$-$00244002\t& WK1002\\_14S &  10:17:57.444  & $-$00:24:40.020  &  1.305  &  2010/02  & 5400 & 6 $\\times$ 900 &  Hn1  & 50  \\\\ S15J142538641+00483135\t& WK1002\\_18S &  14:25:38.639  &  00:48:31.350  &  1.305\t&  2010/02  & 2700 & 3 $\\times$ 900 &  Hn1  & 50  \\\\ S09J090933680+01074587\t& WK1002\\_46S &  09:09:33.680  &  01:07:45.870  &  1.308\t&  2010/02  & 3000 & 5 $\\times$ 600 &  Hn1  & 50  \\\\ S11J110405504+01185565\t& WK1002\\_52S &  11:04:05.504  &  01:18:55.650  &  1.320 \t&  2010/02  & 7200 & 8 $\\times$ 900 &  Hn1  & 50  \\\\ S09J090312056$-$00273273\t& WK0912\\_16S &  09:03:12.056  & $-$00:27:32.730  &  1.327\t&  2009/12\t& 8100 & 9 $\\times$ 900 &  Hn1\t& 50  \\\\  R00J233618953$-$10580749$^{b}$\t& WK0809\\_05R &\t 23:36:18.955 & $-$10:58:07.321     &  1.402\t&  2008/09\t& 2700 & 3 $\\times$ 900 &  Hn2 & 100 \\\\ R00J232248993$-$12151943$^{b}$\t& WK0809\\_04R &  23:22:48.992  & $-$12:15:19.439  &  1.455\t&  2008/09 \t& 2700 & 3 $\\times$ 900 &  Hn3 & 50  \\\\ \\hline \\vspace{0.25mm} \\\\ \\multicolumn{10}{c}{Mis-Identifications}\\\\ \\hline \\vspace{0.25mm} \\\\ R22J214700239+01333423\t& WK0809\\_01R &  21:47:00.239  & 01:33:34.230 &  1.466\t&  09/2008\t& 3600 & 4 $\\times$ 900 &  Hn3\t& 50  \\\\ S00J234420801$-$09322580\t& WK0809\\_03S &  23:44:20.801  &\t$-$09:32:25.800 &  1.488\t&  09/2008\t& 2700 & 3 $\\times$ 900 &  Hn3\t& 50  \\\\ \\hline \\label{osirisobs.table} \\end{tabular*} \\\\ $^{a}$WiggleZ Dark Energy Survey spectroscopic redshift.\\\\ $^{b}$\\halpha detected but observations unusable due to poor AO correction as a result of stray light on the low bandwidth wavefront sensor. \\end{minipage} \\end{table*}   %\\subsection{Disks vs Mergers in IFU Results} A main focus of these studies is the classification of galaxies kinematically to define clear observational distinctions between disks and interacting systems. Mergers observed with integral field units (IFUs; \\citealt{2008A&A...479...67N,2009ApJ...699..421W}) often show clear kinematic steps across one resolution element, indicating separation between two kinematically separate systems, whilst disks are identified by rotational signatures found in both velocity and velocity dispersion maps. A growing number of massive disks have been found with IFS at $z>2$ which are well fit by disk models and which have Toomre parameters of $Q>1$ indicating stable rotationally supported disks, lending observational support to theories of rapid accretion of cold gas and secular disk evolution at high redshift \\citep{2008ApJ...687...59G, 2009A&A...504..789E, 2010MNRAS.402.2291L}.  However, classifying galaxies in these two distinct groups is an over-simplification. Many of the samples present a variety of kinematic types including dispersion-dominated systems, stable disks, unstable disks, minor mergers, mergers, and merger remnants. \\cite{2009ApJ...697.2057L} present a sample of lower mass ($M_{*,avg}=2\\times10^{10}$~\\Msun) galaxies with kinematics dominated by dispersion. Their results can be interpreted as merger remnants, however they also caution that this is statistically unlikely and find that the kinematic classifications are a crude over-simplification. They argue that their galaxies are observed at a range of times during their formation, and that their kinematic signatures are a result of how they acquire their gas. Despite the different kinematic interpretations, all high-redshift IFS samples show high dispersions ($\\sigma \\ga$ 90 km s$^{-1}$) with many (20$-$100\\%) systems showing no ordered rotation making classifications uncertain  (e.g. \\citealt{2008ApJ...687...59G, 2009ApJ...706.1364F, 2009ApJ...697.2057L,2009ApJ...699..421W}).  %\\subsection{Lensed Galaxies} With observation times ranging from 2000s -- 40000s, it is only possible to follow up a small fraction of the brightest galaxies at these redshifts. To supplement high-redshift kinematic results other studies target lensed galaxies magnified in both size and flux, allowing the study of sub-L$_*$ galaxies \\citep{Jones:2010uf}. Using this strategy, IFUs have probed all the way out to $z=4.9$ with NIFS on Gemini \\citep{Swinbank:2009tw} and have targeted strongly lensed galaxies at $z=1.4-3.1$ with OSIRIS on Keck \\citep{2008Natur.455..775S,Jones:2010uf,2011ApJ...732L..14Y}. With galaxy magnification individual clumps have been resolved with diameters of $\\sim$500 pc which will be essential in testing the models of clumpy disk formation \\citep{Jones:2010uf,Swinbank:2009tw,2010Natur.464..733S}. However, more kinematically resolved clumps are needed than are available from current studies ($\\sim 20$ at $z>1$) in order to rigorously compare with models.  Despite great advantages, lensed studies require strong lenses with confirmed redshifts which can be observationally expensive in addition to requiring reliable lensing models and reconstruction. Furthermore, rather than bridge the gap between the $z\\sim2-3$ results and local galaxies, the majority of current lensing studies push out to higher redshifts.  %\\subsection{Previous IFS Studies} In order to correctly interpret IFS data at high redshift it is critical to utilize low-redshift IFS results. By comparing directly to these data and by artificially redshifting them out to $z>1$ a better understanding of different kinematic signatures in the high-redshift data can be reached.  Recent efforts by \\citet{2010MNRAS.401.2113E} have supplied a sizable sample of Fabry\u00d0Perot 3D data of \\halpha emission in 153 nearby isolated galaxies. Disk models are fit to the local population and to the same galaxies artificially redshifted to $z=1.7$, providing guidance in the best disk models to use to fit to high-redshift kinematic results. \\citet{Green:2010fk} and \\citet{2011A&A...527A..60R} also provide local comparison samples for \\halpha velocity dispersion and morphology. Properties observed locally at parsec to kiloparsec resolution are crucial for the interpretation of high-redshift IFU observations.   %\\subsection{THIS PAPER} Connecting the $z<0.2$ and $z>1.5$ results is a relative dearth of IFS studies closer to $z\\sim1$. At this redshift, galaxies are at the tail end of the epoch of mass assembly but at more accessible distances for current instrumentation (where $(1 + z)^{4}$ surface brightness dimming is less severe). Yet suitable galaxies for such studies are rare as once the main epoch of massive galaxy formation is over at $z\\sim2$, the abundance of galaxies with high star formation rates falls rapidly. They are difficult to select with the Lyman-break technique at $z\\sim1$ as the drop in emission blueward of 912\\AA~ cannot be detected with the low sensitivity and resolution of ultraviolet (UV) instruments. These galaxies can, however, be found in very large high-redshift surveys, such as the WiggleZ Dark Energy Survey.  The WiggleZ survey contains a rare population of UV-luminous galaxies out to $z=1.5$ with star formation rates (from nebular emission) of $\\sim$100 M$_{\\odot}$ yr$^{-1}$. These galaxies may be the link needed to bridge the gap between observations of local luminous galaxies, and LBGs and SMGs found at higher redshift. In this paper, we present integral-field spectroscopy of the \\halpha emission line in 13 star-forming galaxies from the WiggleZ Dark Energy Survey at $z\\sim1.3$.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\includegraphics[scale=0.55]{fig1.eps} \\caption{ Nebular luminosity as a function of redshift for the largest high-redshift IFS surveys with the work presented here shown by the red diamonds. The gray points show the \\OII~luminosity of other WiggleZ galaxies considered for this sample. The red circles show the sample from \\citet{2009A&A...504..789E} taken from the VVDS survey and observed with SINFONI on the VLT. The green circles show the sample from \\citet{2009ApJ...697.2057L} observed with OSIRIS on Keck. The blue inverted triangles show the SINS sample from \\citet{2009ApJ...706.1364F} observed with SINFONI, and the green triangles show the sample from \\citet{2009ApJ...699..421W} observed with OSIRIS. The orange points are SMGs with measured \\halpha taken from \\citet{2004ApJ...617...64S} and \\citet{2006ApJ...651..713T}.} \\label{fig.litcomp} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  In this paper, we present 13 galaxies and their physical properties. In Section 2 we describe our sample selection, observational technique, and data reduction methods. In Section 3 we discuss the properties of the galaxies, in particular we describe the morphological and kinematic properties. In Section 3.6 and 3.7 we discuss the methods of disk fitting and surface brightness profile fitting analysis. The individual and global properties of the 13 systems described in Section 3 are analyzed in Section 4 in the context of theories of galaxy formation at high redshift and the implications for the mechanism triggering massive rates of star formation. We conclude with a summary of our results and interpretations in Section 5, giving a brief list of the key points presented in this paper. The properties of the 13 individual galaxies are described individually in the Appendix.  A standard $\\Lambda$CDM cosmology of $\\Omega_{\\mathrm{m}} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $h = 0.7$ is adopted throughout this paper (e.g. \\citealt{2011ApJS..192...18K}). In this cosmology, at redshift $z=1.3$, 1 arcsec corresponds to 8.38 kpc.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We favor the clumpy disk model over the merger sequence model. We find 7/13 of the sample is well modeled by a disk and we find that clumps found in galaxies with multiple regions of emission do not strongly affect the kinematic fields. This provides observational evidence for the models of \\cite{Elmegreen:2008fk} which provide an explanation for the irregular morphologies seen at high redshift.  However, whilst natural to attempt to identify kinematic observations as disks or mergers, as been attempted here, this simplified picture of galaxies becomes increasingly fallible at high-redshift. More than half of the galaxy kinematics in this sample are well-described by disk models, however it is important to be cautious of definitive kinematic classifications. For example, \\cite{2010MNRAS.401.2113E} provide evidence that up to 30\\% of disk galaxies may be misclassified as perturbed rotators or complex kinematics with current classification schemes.  It is likely that we are seeing a variety of mechanisms at work which cannot be broken down into simply disks and mergers. Even if disk instability is a dominant mechanism in star-formation histories at high redshift, as we find here, the disks are still likely to undergo mergers in their lifetime further complicating kinematic interpretations.  \\cite{2010MNRAS.401.1657L}, for example, observe a galaxy pair with separation of 12 kpc with one galaxy in the pair classified as a disk with velocity gradient of $\\sim$91 km s$^{-1}$ and the other galaxy classified is compact and likely dominated by AGN. It is unclear how to classify this system, which has multiple kinematic components. Furthermore, With such small number statistics, the varying kinematics across multiple IFS samples could be due to varying properties amongst the galaxies selected and instruments used, making it difficult to compare results.  The global properties of the sample, the most luminous population, studied at $z\\sim1.3$ with IFS, are in surprising agreement with a model of unstable disks. We agree with \\cite{2009ApJ...697.2057L} that statistically all high-redshift galaxies cannot be mergers. Models show that on average the enhancement of SFR in a merger is only a factor of a few and lasts $200-400$ Myr with peak enhancement lasting for less than 100 Myr \\citep{2010ApJ...720L.149T,2008A&A...492...31D}. It has been shown locally that all ultra-luminous infrared galaxies (ULIRGS) are mergers \\citep{1997A&AS..124..533D} and it is therefore possible a high-redshift subset of galaxies are also exclusively mergers.  However, galaxies on average formed their stars more rapidly in the past implying the need for an evolving definition of what makes more luminous galaxies `extreme'. We therefore stress the importance of the bolometric luminosity in knowing if the sample presented here fall on or above the main sequence of star formation at this redshift to aid in making generalizations about galaxy populations at $z>1$.  %\\subsection{Future Work} %\\subsubsection{Dust} Discerning the dust content will aid in a better understanding of their star-formation histories. Such moderate to high mass, gas rich, intensely luminous galaxies are reminiscent of local ULIRGs and $z\\sim2-3$ SMGs. However, critical to making this link is to determine their far-infrared and sub-mm SEDs, and enabling the measurement of their bolometric luminosities and dust content. Such quantities will inform on whether these galaxies are dust-enshrouded starbursts or \u00d2dust-free\u00d3 late time proto-galaxies. We will present $Herschel$ data in a future paper that will address the dust content in this sample.  %\\subsubsection{Winds} We expect to see signatures of outflows in these systems as galactic scale winds are expected in galaxies forming stars at $>5$\\sfrunits and are seen in 75\\% of ULIRGS \\citep {2005ARA&A..43..769V}. Previous studies have detected galactic winds with IFS (high-z: \\citealt{Swinbank:2009tw,Alexander:2010bd, Shapiro:2009sj}; local: \\citealt{2010ApJ...711..818S}). The high dispersions to the edges of 02S, 52S, and 02R are real and not an inflation from lower signal-to-noise data. This increase in dispersion towards the edges most likely indicates the presence of galactic winds which add an underlying broad component of \\halpha \\citep{2005ARA&A..43..769V}. The influence of winds on this sample will be further discussed in a subsequent paper.  %\\subsubsection{CLUMPS} Measuring the properties of the individual clumps including size, metallicity, SFR, and velocity dispersion to compare to local samples will be important to testing models of clumpy disk formation. Only on order $\\sim10$ clumps have been resolved with IFS studies using lensing and a further $\\sim10$ resolved or marginally resolved within the SINS survey \\citep{2010arXiv1011.5360G}. The properties of the 10 clumps identified in these systems will be presented in our next paper. \\\\  In this paper we:\\\\  \\begin{squishitemize}  \\item Present kinematic data of 13 star-forming galaxies at $z\\sim1.3$, selected from the WiggleZ Dark Energy Survey, of better or comparable resolution than previous studies.  \\item Extend the trend of \\cite{Green:2010fk} to higher star formation rates and higher dispersions between $1<z<2$ (Figure~\\ref{fig.andy}).  \\item Find the most luminous star-forming galaxies at this redshift show a variety of star-forming morphologies with a high incidence of star formation in disks.  \\item Identify clumps of 1--2 kpc in size at $z\\sim1.5$ that are resolved kinematically with LGSAO IFS without aid of lensing.  \\item Conclude the clumps in multiple emission galaxies follow the overall kinematics of the system, and find a high incidence of velocity gradients consistent with rotation lending support of the clumpy disk scenario of EBE08 over the major merger scenario.  \\end{squishitemize}     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1107/1107.2213_arXiv.txt": {
        "abstract": "The topic of this article is weak cosmic shear tomography where the line of sight-weighting is carried out with a set of specifically constructed orthogonal polynomials, dubbed TaRDiS (Tomography with orthogonAl Radial Distance polynomIal Systems). We investigate the properties of these polynomials and employ weak convergence spectra, which have been obtained by weighting with these polynomials, for the estimation of cosmological parameters. We quantify their power in constraining parameters in a Fisher-matrix technique and demonstrate how each polynomial projects out statistically independent information, and how the combination of multiple polynomials lifts degeneracies. The assumption of a reference cosmology is needed for the construction of the polynomials, and as a last point we investigate how errors in the construction with a wrong cosmological model propagate to misestimates in cosmological parameters. TaRDiS performs on a similar level as traditional tomographic methods and some key features of tomography are made easier to understand. ",
        "introduction": "Weak gravitational lensing by the cosmic large-scale structure \\citep{1991MNRAS.251..600B, 1992grle.book.....S, 1994CQGra..11.2345S, 1994A&A...287..349S, 1998MNRAS.301.1064K} is regarded as a very promising way for investigating the properties of the cosmic matter distribution and the measurement of cosmological parameters. The primary tool is the angular spectra of the weak lensing convergence \\citep{1997ApJ...484..560J, 1999ApJ...514L..65H, 2001ApJ...554...67H, 2004PhRvD..70d3009H}, which has, in the context of dark energy cosmologies with adiabatic fluctuations in the dark matter distribution, the potential to deliver parameter constraints competitive with those from the cosmic microwave background.  The weak cosmic shearing effect of the large-scale structure has been detected by four independent research groups a decade ago \\citep{2000A&A...358...30V, 2000astro.ph..3338K, 2000MNRAS.318..625B, 2000Natur.405..143W} and now parameter constraints derived from weak lensing spectra coincide well with those from the CMB, with a small tension concerning the parameters $\\Omega_m$ and $\\sigma_8$. \\citep{2009A&A...497..677K, 2010MNRAS.405.2381K}. These topics are reviewed in detail by \\citet{2001PhR...340..291B} and \\citet{2010RvMP...82..331B, 2010CQGra..27w3001B}.  Together with the property of weak lensing to map out fluctuations in the cosmic matter distribution in a linear way, tomographic methods have been shown to offer superior precision in the determination of cosmological parameters \\citep{1999ApJ...522L..21H, 2002PhRvD..66h3515H, 2002PhRvD..65b3003H, 2004ApJ...601L...1T, 2006JCAP...06..025H}, in particular concerning the properties of dark energy \\citep{2001PhRvD..64l3527H, 2002PhRvD..65f3001H, 2009JCAP...04..012H, 2010GReGr..42.2177H, 2011arXiv1105.0993L, 2011MNRAS.413.1505A}. The idea of tomography is a division of the galaxy sample used for shear estimation into a number of bins in distance. This splitting is the reason for two important advantages of tomography: The weak lensing convergence is a line of sight-integrated quantity of which one measures the angular fluctuation statistics. Angular convergence spectra, however, convey less information compared to the fluctuations statistics of the three-dimensional matter distribution because in the projection process, a mixing of scales is taking place. Secondly, cosmological parameters can have different influences on the amplitude of the weak lensing convergence at different redshifts, which would be averaged out. Combining  convergence spectra measured for each subset of galaxies are able to alleviate these these two problems. Related developments include shear ratio measurements, where the same tidal fields are measured with galaxy populations at different redshifts \\citep{2003PhRvL..91n1302J, 2007MNRAS.374.1377T}, cross correlation cosmography, where part of the galaxies are used for constructing a template on which one observes the shearing effect of more distant galaxies \\citep{2004ApJ...600...17B}, and finally three-dimensional weak shear methods, which provide a direct reconstruction of the fluctuations of the cosmic density field \\citep{2005PhRvD..72b3516C, 2006MNRAS.373..105H, 2011MNRAS.413.2923K}. Common to these methods are nonzero covariances between measurements of different redshifts. This is a natural consequence that lensing is caused by the large-scale structure between the lensed objects and the observer, so that the light from different galaxy samples has to transverse partially the same large-scale structure for reaching the observer.  Our motivation is to revisit weak lensing tomography and to construct a weighing scheme for the galaxies as a function on distance, such that the covariances between spectra estimated from differently weighted galaxy samples has a particular simple, diagonal shape. For this goal, we construct a set of orthogonal polynomials in distance, which provides a diagonalisation of the covariance, and which, when employed in the derivation of the convergence spectrum, projects out statistically independent information about the large-scale structure.  After a brief summary of key formul{\\ae} for cosmology, structure formation and weak lensing in Sect.~\\ref{sect_cosmology}, we describe the construction of TaRDiS-polynomials and investigate their properties in Sect.~\\ref{sect_tomography}. Their statistical and systematical errors are evaluated in Sects.~\\ref{sect_statistics} and~\\ref{sect_systematics}, respectively, and we summarise our results in Sect.~\\ref{sect_summary}. The reference cosmological model used is a spatially flat $w$CDM cosmology with Gaussian adiabatic initial perturbations in the cold dark matter density field. The specific parameter choices are $\\Omega_m = 0.25$, $\\sigma_8 = 0.8$, $H_0=100\\: h\\:\\mathrm{km}/\\mathrm{s}/\\mathrm{Mpc}$, with $h=0.72$, $n_s = 1$ and $\\Omega_b=0.04$. The dark energy equation of state is set to $w=-0.9$ and the sound speed is equal to the speed of light ($c_s=c$) such that there is no clustering in the dark energy fluid. ",
        "conclusions": "\\label{sect_summary} The topic of this paper is a novel method of carrying out tomographic weak lensing measurements by line of sight-weighting of the weak lensing convergence with specifically constructed orthogonal polynomials. \\begin{enumerate} \\item{The TaRDiS-polynomials have been constructed with a Gram-Schmidt orthogonalisation procedure in order to diagonalise the signal covariance matrix, and to project out statistically independent information of the weak lensing field by performing weak lensing tomography. They differ from traditional tomography and 3-dimensional weak lensing methods in the respect that the signal covariance is diagonal instead of the noise covariance.} \\item{The statistical error forcasts for cosmological parameters was investigated with the Fisher-matrix formalism. Because of the polynomial's property of providing statistically independent information of the weak convergence field, one sees a particularly simple square-root scaling of the signal to noise-ratio and the statistical errors with the maximum multipole considered and the number of polynomials used. Clearly, statistical errors can be reduced when combining polynomials, where a small advantage can be observed in comparison to traditional tomography. This is due to the fact that in case of a matching between the assumed and the true cosmology, the Fisher-matrix assumes the largest possible values because of teh diagonalised covariance and provides consequently the smallest statistical errors.} \\item{We extended the Fisher-matrix formalism for investigating how the assumption of a wrong cosmology in the construction of the polynomials impacts on the estimation of cosmological parameters and to what extend biases are introduced. We assumed three systematically wrong priors: assumtion of a time varying instead of a constant dark energy equation of state parameter, as an example of a degree of freedom not contained in the model, a wrong $\\Omega_m$- and $\\sigma_8$-pair, as being the two most prominent parameters determining the strength of the weak lensing signal, and finally a convolved galaxy redshift distribution. All three systematics had a minor impact on the parameter estimation, with the biases decreasing if a larger number of polynomials was used. With 10 polynomials biases were of the order of the statistical error even for strong mismatches in the choice of the initial cosmology. Only the assumption of a wrong redshift distribution generates roughly constant biases, which, when normalised to the statistical errors, increase with the number of polynomials used. Our extension of the Fisher-matrix formalism treats the parameter estimation biases in full generality and does not assume a diagonal shape of neither the signal nor the noise covariance.} \\item{The accuracy needed for constructing viable polynomials corresponds to the statistical error reachable with a simple unweighted convergence spectrum, which is readily available, and can be improved by adding additional priors in the form of the CMB-likelihood, parameter constraints form baryon acoustic oscillations or from supernovae.} \\item{Additionally, it is possible to compute diagnostics for the choice of the prior cosmological model used for constructing the cosmology: Examples include the cross-spectra $S_{ij}(\\ell)$ for two different polynomials $i\\neq j$, which should vanish for correctly-constructed polynomials, and the commutator $\\left[S,N\\right] = SN-NS$ between the signal and the noise covariance, which should vanish if the signal covariance is equal to the unit matrix for a certain choice of normalisation for the polynomials.} \\end{enumerate}  We plan to extend our research on the usage of orthogonal polynomials to the weak lensing bispectrum in a future paper, and to carry out forecasts on dark energy cosmologies from combined constraints with spectrum and bispectrum tomography."
    },
    "1107/1107.2025_arXiv.txt": {
        "abstract": "\\centerline{\\bf Abstract} We study the evolution of primordail black holes by considering present universe is no more matter dominated rather vacuum energy dominated. We also consider the accretion of radiation, matter and vacuum energy during respective dominance period. In this scenario, we found that radiation accretion efficiency should be less than $0.366$ and accretion rate is much larger than previous analysis by B. Nayak et al. \\cite{ns}. Thus here primordial black holes live longer than previous works \\cite{ns}. Again matter accretion slightly increases the mass and lifetime of primordial black holes. However, the vacuum energy accretion is slightly complicated one, where accretion is possible only upto a critical time. This critical time depends on the values of accretion efficiency and formation time. If a primordial black hole lives beyond critical time, then its' lifespan increases due to  vacuum energy accretion. But for presently evaporating primordial black holes, critical time comes much later than their evaporating time and  thus vacuum energy could not affect those primordial black holes. We again found that the constraints on the initial mass fraction of PBH obtained from the $\\gamma$-ray background limit becomes stronger in the presence of vacuum energy. ",
        "introduction": "Black holes which are formed in the early universe are known as Primordial Black Holes (PBHs). A comparison of the cosmological density of the universe at any time after the Big Bang with the density associated with a black hole shows that PBHs would have of order the particle horizon mass. PBHs could thus span enormous mass range starting from $10^{-5}gm$ to more than $10^{15}gm$. These black holes are formed as a result of initial inhomogeneities \\cite{zn, bjc}, inflation \\cite{kmz, cgl}, phase transitions \\cite{kp}, bubble collisions \\cite{kss, ls}, or the decay of cosmic loops \\cite{pz}. In 1974 Hawking discovered that the black holes emit thermal radiation due to quantum effects \\cite{haw}. So the black holes get evaporated depending upon their masses. But in references \\cite{ns}, it is shown that evaporation of primordial black holes delayed due to accretion of radiation by assuming standard picture of Cosmology. Similar kind of works have been done by many other authors \\cite{mj, fca, jdj}.  In standard picture of Cosmology \\cite{sw}, universe is radiation dominated in early period of evolution and is matter dominated now. So universe is undergoing a decelerated expansion through out its evolution. But observations of distant supernovae and cosmic microwave background anisotropy indicates that the present universe is undergoing accelerating expansion \\cite{apj}. To explain this unwanted observational fact it is thought that present universe is dominated by vacuum energy with negative pressure termed as dark energy. SN Ia observations also provide the evidence of a decelerated universe in the recent past with trasition from decelerated to accelerated occuring at redshift $z_{q=0}\\sim 0.5$ \\cite{mst}. So the vacuum energy domination should be started from $z_{q=0}\\sim0.5$ i.e. $t_{q=0}\\sim\\frac{1}{2}t_0$.  In this work, we study the evolution of PBH by considering present universe is vacuum energy dominated. Here we consider accretion of radiation, matter and vacuum energy during respective dominance period and mainly discuss how accretion of vacuum energy affect PBH evolution. ",
        "conclusions": "Here we study the evolution of primordail black holes by considering present universe is vacuum energy dominated. In our consideration, we have taken that the universe evolves from an initial radiation dominated to matter dominated and finally to present vacuum energy phase. We also consider that when a PBH passes through radiation domination, matter domination and vacuum domination, it accretes radiation, matter and vacuum energy respectively. During radiation dominated era, we found that radiation accretion efficiency should less than $0.366$ and accretion rate is much larger than previous works by B. Nayak et al. \\cite{ns}. Thus primordial black holes live longer than previous analysis \\cite{ns}. In matter dominated era, accretion of matter slightly increases the mass and lifetime of primordial black holes. However, during vacuum energy dominated era, the accumulation of vacuum energy is possible only upto a critical time $t_c$. The value of $t_c$ depends on accretion efficiency and formation time. If a PBH lives beyond this critical time, then its' lifespan increases due to accretion of vacuum energy. But for presently evaporating PBHs, the critial time $t_c$ comes much later than their evaporating time. So those PBHs are not affected by the presence of vacuum energy. We also found that the constraint on the initial mass fraction of PBH obtained from the $\\gamma-$ray background limit becomes stronger in the presence of vacuum energy."
    },
    "1107/1107.5119_arXiv.txt": {
        "abstract": "{The main goal of our project is to obtain a complete picture of individual open clusters from homogeneous data and then search for correlations between their astrophysical parameters. The near--infrared $JHK_{S}$ photometric data from the 2--Micron All Sky Survey were used to determine new coordinates of the centres, angular sizes and radial density profiles for 849 open clusters in the Milky Way. Additionally, age, reddening, distance, and linear sizes were also derived for 754 of them. For \\textbf{these} open clusters our results are in satisfactory agreement with the literature data. The analysed sample contains open clusters with ages in the range from 7 Myr to 10 Gyr. The majority of these clusters are located up to 3 kpc from the Sun, less than 0.4 kpc from the Galactic Plane and 6 -- 12 kpc from the Galactic Centre. The majority of clusters have core radii of about 1.5 pc and the limiting radii of the order of 10 pc. We notice that in the near--infrared, open clusters seem to be greater than in optical bands. We notice that a paucity of clusters is observed at Galactic longitudes range from $140\\,^{\\circ}$ to $200\\,^{\\circ}$ which probably reflects the real spatial distribution of open clusters in the Galaxy. The lack of clusters was also found in earlier studies.}  {open clusters and associations: general -- infrared: galaxies -- astronomical databases: 2MASS} ",
        "introduction": "In this paper the near-infrared $JHK_{S}$ photometric data were used to redetermine equatorial and galactic coordinates of the centres of known open clusters. This allowed us to construct radial density profiles of clusters and to determine their core radii and estimate their limiting radii. The adopted background-star decontaminating procedure allowed us to construct colour--magnitude diagrams for individual clusters and to estimate ages for a significant fraction of the considered sample.  The structure of this paper is as follows. In Section 2 the method and data analysis is presented. In Section 3 our results are compared to the results found in literature. Section 4 contains discussion of the relations between individual  parameters. Final conclusions are summarized in Section 5.  \\textbf{Similar to our research, Janes (1979) used $\\mathitbf{UBV}$ photometry to study the reddening and metallicity of 41 open clusters. Comparison between the age and position in the Galaxy was studied by Lyng\\aa~(1980, 1982). Janes \\& Phelps (1994) based on 72 open clusters and Friel (1995) with \\textbf{a} sample \\textbf{of} 74 objects \\textbf{investigated} how the extinction and age depend on \\textbf{a} position in the Galaxy. Tadross et al.~(2002) and Schilbach et al.~(2006) studied how open clusters diameters and age depend on position in the Galactic Plane. Previous research of open clusters in 2MASS photometry were done by Dutra \\& Bica (2001) who studied 42 new infrared star clusters, stellar groups and candidates towards the Cyngus X region. Bica et al.~(2003) based on 2MASS \\textbf{researched} 346 open clusters, whereas 315 objects were previously known. They studied linear diameters and spatial distribution of open clusters in the Galaxy. Recently Froebrich et al.~(2010) studied 269 open clusters and the age, core radius, the reddening, Galactocentric distance and the scaleheight were determined.} ",
        "conclusions": "The near--infrared $JHK_{S}$ photometric data from the 2MASS \\emph{Point Source Catalogue} (Skrutskie et al.~ 2006) was used to determine structural and physical parameters of a large sample of known open clusters. The structural properties have been determined for the complete sample of 849 open clusters, while the age and distance were obtained for 754 of them.  We showed that open clusters studied in near-infrared seem to be larger than in optical bands. The same property was noticed by Sharma et al.~ (2006). Near IR data allow us to detect faint cluster members located far from the cluster core and are not limited by the telescope's field of view.  We obtained a different relation of $E(J - K_{S})/ E(B - V)$ than the one given by Schlegel et al.~(1998) who obtained $E(J - K_{S})/ E(B - V) = \\mathbf{0.520}$. Our results give $E(J - K_{S})/ E(B - V) = \\mathbf{0.463}$. This difference translates into the error of the distance determinations of only \\textbf{2.8}{\\%}. We noticed that when moving outward from the Galactic Plane and the Galactic Centre the reddening of clusters statistically decreases. The mean $E(B - V)$ for $|Z| < 0.1$ kpc is \\textbf{0.68} mag, between 0.1 and 0.4 kpc is equal to \\textbf{0.51} mag and for $|Z| > 0.4$ kpc is only \\textbf{0.34} mag. We showed the similar trend for $R_{GC}$. For clusters located near the Sun's orbit and further than 8.3 kpc from the Galactic Centre, the mean $E(B - V)$ is \\textbf{0.74} and \\textbf{0.53} mag, respectively. This may be implied by the fact that the dust and gas density decreases, too.  For the majority of studied clusters (81{\\%}) the limiting radii was found to be smaller than 10 pc and the core radii smaller than 1.5 pc. Statistically the limiting radius is about 6--7 times larger than the core radius. We also noticed that the average cluster sizes appear to grow with a distance from the Galactic Centre and Galactic Plane. This effect has already been observed by Lyng\\aa~(1982), van den Bergh et al.~ (1991), Tadross et al.~(2002), Schilbach et al.~(2006) and Bonatto and Bica (2010). Moreover, we noticed that the average age of clusters increases with a distance from the Galactic Centre and the Galactic Plane what was shown by Janes and Phelps (1994), Friel (1995), Tadross et al.~(2002) and Froebrich et al.~(2010).  Open clusters, younger than 500 Myr, have a scale height of 71 pc, while older clusters 238 pc. These results are similar to Janes and Phelps (1994) who derived scale heights of 55 pc and 375 pc for young and old clusters, respectively. We found that a number of older clusters increases with the distance from the Galactic Centre. For younger clusters we obtained an average $R_{GC} = 9.0$ kpc and for the older ones -- $10.0$ kpc. This seems to confirm values of $8.9$ and $9.4$ kpc for younger and older clusters, respectively, obtained by Froebrich et al.~(2010) and by Tadross et al.~(2002) who found that younger clusters are concentrated at $R_{GC} < 9.5$ kpc. This shows a difference between younger and older clusters in their location. We also notice that older clusters seem to be larger than the younger ones, what was shown by Lyng\\aa~(1982) and Tadross et al.~(2002).  Moreover we found important selection effects in our sample. Cluster sizes grow with distance from the Sun, the same problem was encountered in Froebrich et al.~(2010).  We observed a smaller number of clusters at the Galactic longitude range of $140\\,^{\\circ} < l < 200\\,^{\\circ}$ what is in agreement with Tadross et al.~(2002) and Froebrich et al.~(2010). This may reflect the structure of the Galaxy in this direction where there is a gap in the Perseus Arm (Benjamin 2008).  \\Acknow{This research is supported by \"Stypendia dla doktorant\\'ow 2008/2009 -- ZPORR\" SPS.IV-3040-UE/204/2009. 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.}"
    },
    "1107/1107.2952_arXiv.txt": {
        "abstract": "We present an analysis of the extended emission around the anomalous X-ray pulsar 1E 1547.0$-$5408 using four \\textit{XMM-Newton} observations taken with the source in varying states of outburst as well as in quiescence. We find that the extended emission flux is highly variable and strongly correlated with the flux of the magnetar. Based on this result, as well as on spectral and energetic considerations, we conclude that the extended emission is dominated by a dust-scattering halo and not a pulsar wind nebula (PWN), as has been previously argued. We obtain an upper limit on the 2--10\\,keV flux of a possible PWN of $4.7\\times10^{-14}\\,\\mathrm{erg\\, s^{-1}\\, cm^{-2}}$, three times less than the previously claimed value, implying an efficiency for conversion of spin-down energy into nebular luminosity of $<$$9\\times10^{-4}$ (assuming a distance of 4\\,kpc). We do, however, find strong evidence for X-ray emission from the supernova remnant shell surrounding the pulsar, as previously reported. ",
        "introduction": "Anomalous X-ray pulsars (AXPs) and soft gamma repeaters (SGRs) are generally accepted as belonging to the class of neutron stars known as `magnetars.' These sources are characterized by long (2--12\\,s) spin periods and spin-down rates that imply extremely high surface dipole magnetic fields of $10^{14}\\textrm{--}10^{15}\\,\\mathrm{G}$ (see \\citealt{wt06} and \\citealt{m08} for reviews). It is thought that their large X-ray luminosities are powered by the decay of these powerful magnetic fields, and fracturing of the crust and reconfiguration of the magnetic field lines produce the glitches, bursts, and X-ray variability seen in these sources \\citep{td95,td96,tlk02}.  One important open question regarding AXPs and SGRs is whether they produce particle outflows akin to those seen in conventional rotation-powered pulsars. The latter are well known to produce often spectacular `pulsar wind nebulae' (PWNe), the classic example of which is the Crab Nebula \\citep[see, e.g.,][for reviews]{krh06,gs06,kp08}. Such nebulae, evident particularly at radio and X-ray energies, are the result of synchrotron emission due to pulsar-produced relativistic electrons and positrons, as they spiral in the ambient magnetic field. Given the large magnetic energy reservoir hypothesized to exist in AXPs and SGRs, particle outflows seem reasonable to consider, and indeed have been suggested to be present ubiquitously in magnetars, either in continuous or sporadic forms \\citep{tb98,hck99}. This idea was initially buoyed by the claim of an apparent wind nebula associated with SGR 1806$-$20 \\citep{mtk+94}. Although the latter association was later disproved \\citep{hkc+99}, the possibility of a nebula-producing magnetar wind has not been. Extended radio emission was unambiguously identified following one SGR giant flare \\citep{gkg+05,tgg+05,gle+05,fmg+06}, but is thought to be from relativistic, weakly baryon-loaded magnetic clouds \\citep{l06} or from a baryonic outflow \\citep{gle+05,g07}, and associated exclusively with the flare. Recently \\citet{rmg+09} have suggested that an unusual X-ray nebula surrounding the relatively high-$B$ Rotating Radio Transient (RRAT) J1819$-$1458 is magnetically powered, though the mechanism for this is unclear.  The X-ray source 1E 1547.0$-$5408 was discovered with the \\textit{Einstein} X-ray satellite in 1980 by \\citet{lm81}, but only recently was it suggested to be a magnetar candidate on the basis of its spectrum, variable nature, and likely association with the supernova remnant (SNR) shell G327.24--0.13 \\citep{gg07}. Radio observations of the source by \\citet{crhr07} revealed pulsations at a period of 2.1\\,s and the measured spin-down rate implied a surface magnetic field strength of $2.2\\times10^{14}\\,\\mathrm{G}$ and a spin-down luminosity of $\\dot{E}=1.0\\times10^{35}\\,\\mathrm{erg\\, s^{-1}}$. 1E 1547.0$-$5408 is thus the fastest-rotating and highest-$\\dot{E}$ magnetar yet known.% \\footnote{See the McGill SGR/AXP Online Catalog at \\url{http://www.physics.mcgill.ca/~pulsar/magnetar/main.html}.% } Although observations in 2006 showed the source to be in quiescence, a 2007 \\textit{XMM-Newton} observation showed it to be in a high, apparently post-outburst state \\citep{hgr+08}. In 2008 October and again on 2009 January 29, 1E 1547.0$-$5408 underwent strong outburst events, experiencing dramatic increases in its X-ray luminosity and exhibiting many SGR-like bursts within a few hours. For further details on these outbursts and the source's history, see \\citet{ier+10}, \\citet{kgk+10}, \\citet{nkd+11}, \\citet{bis+11}, \\citet{sk11}, and \\citet{dkg11}.  In observations of 1E 1547.0$-$5408 taken with \\textit{Swift} and \\textit{XMM-Newton} following the 2009 outburst, \\citet{tve+10} observed dust-scattering X-ray rings centered on the magnetar and derived from them a source distance of $\\sim$4--5\\,kpc. They also found evidence of time variable diffuse emission around the source which they attributed to dust scattering of the bursts and of the persistant X-ray emission.  Meanwhile, \\citet[hereafter VB09]{vb09}, analyzing \\textit{Chandra} and \\textit{XMM-Newton} observations of the source taken in 2006 when it was in quiescence, detected extended emission and characterized it as the result of a PWN, in analogy with those seen around rotation-powered pulsars. They argued for a PWN based on the high flux level of the extended emission, and because it appeared to have a harder spectrum than the point source. This would make 1E 1547.0$-$5408 unusual among the known magnetars, as no other such source has been shown to power such emission. They also showed the presence of extended X-ray emission coincident with the SNR shell.  Here we present an analysis of the extended emission around 1E 1547.0$-$5408 using multi-epoch \\textit{XMM} data, in which the source flux varies strongly, as does the putative nebular emission. We show conclusively that the putative PWN reported by VB09 is in fact dominated by dust scattering, rather than by emission from any pulsar outflow. ",
        "conclusions": "In this paper, we have examined multi-epoch \\textit{XMM-Newton} data for the magnetar candidate 1E 1547.0$-$5408 and we show that the observed extended emission surrounding the source is dominated by dust-scattered magnetar emission. Specifically we find that the luminosity of the nebular emission is proportional to the source flux, as expected for dust scattering, but not seen in any known PWN or other magnetar candidate. Additional strong evidence for dust-scattering comes from spectral and energetics arguments, as well as from the disagreement between the observed nebular variability time scale and the expected synchrotron loss time in the PWN interpretation. We note that contrary to a previous claim (VB09), even in 2006 when the source was relatively faint, $64\\%\\pm12\\%$ of the nebular emission is from dust scattering. We cannot, however, rule out the presence of a faint PWN with luminosity $\\lesssim$$9\\times10^{31}(d/4\\,\\mathrm{kpc})^{2}\\,\\mathrm{erg\\, s^{-1}}$ in the 2--10 keV band, a limit three times lower than the previously claimed detection ($\\sim$$2.9\\times10^{32}\\,\\mathrm{erg\\, s^{-1}}$ from VB09). Deep observations of this source when the magnetar is in quiescence are necessary to test this hypothesis. We do, on the other hand, find strong evidence for non-dust-scattered extended X-ray emission at angular distance $\\sim$$40\\arcsec\\textrm{--}150\\arcsec$, which we argue is from the SNR shell surrounding the pulsar, as previously reported by VB09.  With the absence of evidence for a PWN surrounding any AXP or SGR, now including 1E 1547.0$-$5408, previous models for the production of PWNe by magnetars \\citep{tb98,hck99} remain unsupported. On the other hand with clear evidence for the existence of PWNe surrounding many high-magnetic-field radio pulsars (e.g., PSR J1846--0258: \\citealp{hcg03}; PSR J1119--6127: \\citealp{gs03}) which generally have substantially higher spin-down luminosities than any known AXP or SGR, the production of the relativistic particle wind necessary to generate an observable PWN seems intimately tied to the rotation-derived power, rather than that from magnetic-field decay. This then makes the detection of the surprisingly bright X-ray PWN surrounding the presumably rotation-powered but relatively high-$B$ RRAT PSR J1819$-$1458 \\citep{rmg+09} particularly interesting and worthy of follow-up."
    },
    "1107/1107.5096_arXiv.txt": {
        "abstract": "{% We present a near-infrared extinction map of a large region (approximately $2\\,200 \\mbox{ deg}^2$) covering the Orion, the Monoceros R2, the Rosette, and the Canis Major molecular clouds.  We used robust and optimal methods to map the dust column density in the near-infrared (\\textsc{Nicer} and \\textsc{Nicest}) towards $\\sim 19$ million stars of the Two Micron All Sky Survey (2MASS) point source catalog.  Over the relevant regions of the field, we reached a 1-$\\sigma$ error of $0.03 \\mbox{ mag}$ in the $K$-band extinction with a resolution of $3 \\mbox{ arcmin}$.  We measured the cloud distances by comparing the observed density of foreground stars with the prediction of galactic models, thus obtaining $d_\\mathrm{Orion A} = (371 \\pm 10) \\mbox{ pc}$, $d_\\mathrm{Orion B} = (398 \\pm 12) \\mbox{ pc}$, $d_\\mathrm{Mon R2} = (905 \\pm 37) \\mbox{ pc}$, $d_\\mathrm{Rosette} = (1330 \\pm 48) \\mbox{ pc}$, and $d_\\mathrm{CMa} = (1150 \\pm 64) \\mbox{ pc}$, values that compare very well with independent estimates. ",
        "introduction": "\\label{sec:introduction}  In a series of papers, we have applied an optimized multi-band technique dubbed Near-Infrared Color Excess Revisited (\\textsc{Nicer} \\citealp{2001A&A...377.1023L}, hereafter Paper~0) to measure dust extinction and investigate the structure of nearby molecular dark clouds using the Two Micron All Sky Survey (2MASS; \\citealp{1994ExA.....3...65K}).  The main aim of our coordinated study is to investigate the large-scale structure of these objects and to clarify the link between the global physical properties of molecular clouds and their ability to form stars.  Previously, we considered the Pipe nebula (see \\citealp{2006A&A...454..781L}, hereafter Paper~I), the Ophiuchus and Lupus complexes (\\citealp{2008A&A...489..143L}, hereafter Paper~II), and the Taurus, Perseus, and California complexes (\\citealp{2010A&A...512A..67L}, hereafter Paper~III).  We now present an analysis of a large region covering more than $2\\,200$ square degrees, centered around Orion.  This region includes Orion~A and B, $\\lambda$ Orionis, Mon R2, Rosette, and Canis Major.  An overview of a subset of the wide field extinction map presented in this paper, superimposed on an optical image of the sky, is presented in Fig.~\\ref{fig:1}.  \\begin{figure*}[!tbp] \\centering \\includegraphics[width=15cm]{fig01a.jpg}% \\hspace{-15cm}% \\begin{ocg}{clouds}{1}{1}% \\includegraphics[width=15cm]{fig01b.jpg}% \\end{ocg}% \\hspace{-15cm}% \\begin{ocg}{labels}{2}{1}% \\includegraphics[width=15cm]{fig01c.pdf}% \\end{ocg}% \\caption{An optical image of Orion (Wei-Hao Wang, IfA, University of Hawaii) with the extinction map presented in this paper superimposed in green.  The complementarity between the red \\textsc{Hii} regions and the green $\\mathrm{H}_2$ ones is evident. Note that the online figure is interactive: \\ToggleLayer{clouds}{\\protect\\framebox{extinction}} and \\ToggleLayer{labels}{\\protect\\framebox{labels}} can be toggled by clicking on the respective boxes when using Adobe$\\textsuperscript{\\textregistered}$ Acrobat$\\textsuperscript{\\textregistered}$.} \\label{fig:1} \\end{figure*}  Near-infrared dust extinction measurement techniques present several advantages with respect to other column density tracers.  As shown by \\citet{2009ApJ...692...91G}, observations of dust are a better column density tracer than observations of molecular gas (CO), and observations of dust extinction in particular provide more robust measurements of column density than observations of dust emission, mainly because of the dependence of the latter measurements on uncertain knowledge of dust temperatures and emissivities. Additionally, the sensitivity reached by near-infrared dust extinction techniques, and by \\textsc{Nicer} in particular, is essential to investigate the low-density regions of molecular clouds (which are often below the column density threshold required for the detection of the CO molecule) and therefore to estimate the mass of the diffuse gas that acts as a pressure boundary around the clumps.  Similar to Paper~III, we use for some key analyses in this paper the improved \\textsc{Nicest} method \\citep{2009A&A...493..735L}, designed to cope better with the unresolved inhomogeneities present in the high-column density regions of the maps.  Orion is probably the best studied molecular cloud in the sky.  The complex comprises \\textsc{Hii} and \\textsc{Hi} regions superimposed on colder, massive H$_2$ clouds with active formation of both low and high mass stars.  The area studied here contains the Orion OB association, which is split in several subgroups with ages from $2$ to $12 \\mbox{ Myr}$.  It is estimated that in the last $12 \\mbox{ Myr}$ there have been 10 to 20 supernov\\ae\\ explosions \\citep{2008hsf1.book..459B} that shaped the gas in the region.  More than a century ago, Barnard discovered a large arc of H$_\\alpha$ emission around the eastern part of Orion (see Fig.~\\ref{fig:1}). More recently, Barnard's loop has been linked to Eridanus loop, and the two have been identified as a superbubble that is expanding into denser regions with a mean velocity of $10$--$20 \\mbox{ km s}^{-1}$ \\citep{1988ApJ...324..776M}.  Orion includes two giant molecular molecular clouds, Orion~A and Orion~B, the spectacular $\\lambda$ Orionis bubble, and a large number of smaller cometary clouds.  Both Orion~A and B are observed in projection inside Barnard's loop and most likely have been shaped by the supernova explosions, stellar winds, and \\textsc{Hii} regions generated by OB stars in the area; for example, note the shell originating from the south-east extremity of Orion~A, around the early B star $\\kappa$ Orionis (Saiph).  Orion~A and B host rich clusters of young ($\\sim 2 \\mbox{ Myr}$) stars, the Orion Nebula Cluster (ONC), NGC~2024 (also known as the Flame Nebula), NGC~2071, and NGC~2068.  Just southwest of NGC~2024 is the famous Horsehead Nebula clearly visible as a protrusion of the green extinction map in Fig.~\\ref{fig:1}.  Around the ``head'' of Orion, the O8 star $\\lambda$ Orionis, there is a ring of dark clouds, $60 \\mbox{ pc}$ diameter, known for almost a century.  The region hosts several young stars (including 11 OB stars very close to $\\lambda$ Orionis), and their spatial and age distributions show that originally star formation occurred in an elongated giant molecular cloud \\citep{2008hsf1.book..757M, 1982ApJ...261..135D}.  Most likely, the explosion of a supernova coupled with stellar winds and \\textsc{Hii} regions destroyed the dense central core, created an ionized bubble and a molecular shell visible as a ring.  Star formation still continues in remnant dark clouds distant from the original core.  The Monoceros R2 region is distinguished by a chain of reflection nebul\\ae\\ that extend over $2^\\circ$ on the sky.  The nomenclature ``Mon R2'' indicates the second association of reflection nebul\\ae\\ in the constellation Monoceros \\citet{1966AJ.....71..990V}.  The region is also well known for its dark nebula, which is clearly visible on top of the reflection nebul\\ae\\ and the field stars. Clusters of newly formed stars are present in its core and in the GGD 12-15 region; smaller cores are also present in the field.  The cloud is estimated to be at a distance significantly higher than the Orion nebula, $830 \\pm 50 \\mbox{ pc}$ \\citep{1976AJ.....81..840H}.  Rosette is known to amateur astronomers as one of the most spectacular nebul\\ae\\ in the sky.  The region is characterized by an expanding \\textsc{Hii} region interacting with a giant molecular cloud.  The \\textsc{Hii} region has been photodissociated by NGC~2244, a cluster containing more than 30 high-mass OB stars.  Judging from the age of NGC~2244, the Rosette molecular cloud has been actively forming stars for the last 2 to 3 Myr, and most likely will continue doing so, as there is still sufficient molecular material placed in a heavily stimulated environment \\citep{2008hsf1.book..928R}.  The distance to NGC~2244, and therefore to the molecular cloud, is still controversial: recent estimates range between $1\\,390 \\mbox{ pc}$ \\citep{2000A&A...358..553H} and $1\\,670 \\mbox{ pc}$ \\citep{2002AJ....123..892P}.  The star-forming region in Canis Major is characterized by a concentrated group of early type stars (forming the CMa OB1/R1 associations) and by an arc-shaped molecular cloud (probably produced by a supernova explosion).  Similarly to Rosette, Canis Major presents an interface between the \\textsc{Hii} region and the neutral gas which shows up as a north-south oriented ridge.  Recent distance determinations of Canis Major seem to converge around a distance of $1\\,000 \\mbox{ pc}$ \\citep{2008hsf2.book....1G}.  This paper is organized as follows.  In Sect.~\\ref{sec:nicer-absorpt-map} we briefly describe the technique used to map the dust and we present the main results obtained. Section~\\ref{sec:statistical-analysis} is devoted to an in-depth statistical analysis, and includes a discussion of the measured reddening laws for the various clouds, the measurements of the distances using foreground stars, the log-normality of the column density distributions, and the effects of small-scale inhomogeneities. The mass estimates for the clouds are presented in Sect.~\\ref{sec:mass-estimate}.  Finally, we summarize the results obtained in Sect.~\\ref{sec:conclusions}.  A few plots in this paper are associated with switches designed to show or hide labels; these are shown as frames in the respective captions.  In order to use this feature this electronic document should be displayed using Adobe$\\textsuperscript{\\textregistered}$ Acrobat$\\textsuperscript{\\textregistered}$. ",
        "conclusions": "\\label{sec:conclusions}  The following items summarize the main results presented in this paper: \\begin{itemize} \\item We measured the extinction over an area of $\\sim 2\\,200$ square degrees that encompasses the Orion, the Monoceros R2, the Rosette, and the Canis Major molecular clouds.  The extinction map, obtained with the \\textsc{Nicer} and \\textsc{Nicest} algorithms, has a resolution of $3 \\mbox{ arcmin}$ and an average 1-$\\sigma$ detection level of $0.3$ visual magnitudes. \\item We measured the reddening laws for the various clouds and showed that they agree with the standard \\citet{2005ApJ...619..931I} reddening law, with the exception of the Mon~R2 region.  We interpret the discrepancies observed as a result of contamination from a population of young blue stars (likely including members of the Orion OB1 association). \\item We estimated the distances of the various clouds by comparing the density of foreground stars with the prediction of the \\citet{2003A&A...409..523R} Galactic model.  All values obtained were found to be in very good agreement with independent measurements. \\item We considered the column density probability distributions for the clouds and obtained reasonable log-normal fits for all of them. \\item We measured the masses of the clouds and their cumulative mass distributions and found differences in the internal structure among the clouds. \\end{itemize}"
    },
    "1107/1107.5575_arXiv.txt": {
        "abstract": "We study effects of particle re-acceleration (or heating) in the post-shock region via magnetohydrodynamic/plasma turbulence, in the context of a mixed hadronic-leptonic model for the prompt emission of gamma-ray bursts (GRBs), using both analytical and numerical methods.  We show that stochastically accelerated (or heated) leptons, which are injected via $pp$ and $p \\gamma$ reactions and subsequent pair cascades, are plausibly able to reproduce the Band function spectra with $\\alpha \\sim 1$ and $\\beta \\sim 2-3$ in the $\\sim$~MeV range.  An additional hard component coming from the proton-induced cascade emission is simultaneously expected, which is compatible with observed extra power-law spectra far above the MeV range.  We also discuss the specific implications of hadronic models for ongoing high-energy neutrino observations. ",
        "introduction": "The origin of prompt emission from gamma-ray bursts (GRBs) has been an open issue for more than 40 years.  It is mainly observed in the $E\\sim E^{\\rm br} \\sim 10~{\\rm keV}-1~ {\\rm MeV}$ range, most spectra being fitted by a smoothed broken power-law (the so-called Band function) with break energy $E^{\\rm br}$.  The typical low-energy photon index is $\\alpha \\sim 1$ (where $d F/d E \\propto E^{-\\alpha}$ below $\\sim E^{\\rm br}$), while the typical high-energy photon index is $\\beta \\sim 2-3$ (where $d F/d E \\propto E^{-\\beta}$ above $\\sim E^{\\rm br}$).  Observed light curves are highly variable.  The variability can sometimes be $\\sim$~ms, but pulses with various widths are observed.  For typical long GRBs, light curves consist of many variable pulses, so that the duration, $T \\sim {10}^{1-2}$~s, is usually longer.  So far, many models have been proposed to explain this most luminous phenomenon in our universe~\\citep[see recent reviews, e.g.,][]{Mes06,Zha07}.  The classical scenario of the prompt emission is the optically-thin synchrotron internal shock model~\\citep[e.g.,][]{RM94}, where shocks are responsible for dissipation of the outflow kinetic energy and observed gamma rays are attributed to synchrotron radiation from relativistic electrons. However, there are a couple of issues in the classical scenario.  First, the low-energy photon index may be incompatible with values predicted by synchrotron emission~\\cite[e.g.,][]{GC99,MR00}.  This is the case especially when relativistic electrons are in the fast-cooling regime, which leads to $\\alpha \\sim 1.5$, although the Klein-Nishina effect on synchrotron self-Compton (SSC) process may reproduce $\\alpha \\sim 1$~\\citep[e.g.,][]{DKK01,WLDM09,DBD11}.  One of the possible resolutions of this fast-cooling problem is to consider the magnetic field decay timescale that is comparable to the electron cooling timescale~\\cite{PZ06}.  Injection of electrons to acceleration processes is one of the related issues.  Shock acceleration of primary electrons may be inefficient for sufficiently small emission radii because of the rapid cooling in the strong magnetic field region at the vicinity of the shock~\\cite{MS09}.  If the outflow is magnetized enough, even for sub-luminal shocks, recent \\textit{ab intio} particle-in-cell simulations of collisionless shocks indicate that acceleration of nonthermal electrons at relativistic shocks are inefficient whereas ions can be accelerated~\\cite{SS11}. Another issue is the radiative efficiency problem~\\citep[e.g.,][]{ITYN06,Zha+07}, although this depends on estimates of the afterglow kinetic energy~\\cite{EW05}.  The high efficiency can be achieved by a large dispersion in the Lorentz factor of the outflows~\\citep[e.g.,][]{Bel00}, but it does not seem so easy to reconcile that with the observed spectral correlations~\\citep[e.g.,][]{Yon+04}.  A recent discussion of the classical scenario is found in Daigne et al. (2011).  On the other hand, many alternative models have been suggested.  Diffusive synchrotron radiation can explain the low-energy photon index if the fast-cooling problem is fixed~\\cite{Med00}.  Instead of dissipation of the kinetic energy via internal shocks, dissipation of the magnetic energy via, e.g., magnetic reconnection may play a crucial role as the dissipation mechanism, and observed gamma rays may be synchrotron or other electromagnetic radiation from relativistic electrons~\\citep[e.g.,][]{SDD01,Lyu06,ZY11,MU11}. If significant thermal energy is stored in the GRB jet by energy deposition from the central engine and/or via some dissipation below the photosphere, the prompt emission may originate from quasi-thermal emission caused by Comptonization or Coulomb heating~\\citep[e.g.,][]{Tho94,GC99,MR00,PMR06,Iok+07,Bel10}. One of the appealing possibilities is to consider the slow heating scenario.  Around the photosphere, electrons that are slowly heated can lead to $\\alpha \\sim 1$ via Comptonization~\\cite{Tho94,PMR06}, where the typical electron energy is determined by the balance between heating and cooling~\\cite{GC99,VP09}.  Alternatively, in the neutron-rich outflow, Coulomb collisions should work as slow heating below the baryonic photosphere, leading to hard photon spectra of $\\alpha \\sim -0.5-1$ and $\\beta \\sim 2.5$ via the pair injection by the $np$ reaction and subsequent cascades~\\cite{Bel10,VBP11}.  Such slow heating may be caused by stochastic acceleration (so-called second-order Fermi acceleration) over the sub-hydrodynamical timescale and/or hydrodynamical timescale, which may occur in the magnetohydrodynamic (MHD) and plasma turbulences generated by shock or magnetic dissipation such as magnetic reconnections~\\cite{BM96,AT09,ZY11}.  Emission at higher energies above $\\sim 10$~MeV should also be important to reveal the radiation mechanism of the prompt emission, and various high-energy processes such as Comptonized thermal, SSC~\\cite{RM94,AI07}, proton-induced cascade~\\cite{Vie97,DA06}, and proton synchrotron~\\cite{Tot98} emissions, were discussed.  Such high-energy emission had been sparsely detected by the EGRET onboard the Compton Gamma-Ray Observatory~\\citep[e.g.,][]{Hur+94}.  Recently, however, significant observational progress has been achieved by \\textit{Fermi}.  Dozens of GRBs have been observed by LAT onboard \\textit{Fermi}, and some bursts can be fitted by the Band function up to GeV energies~\\cite{Abd+09a,Abd+09c,Abd+10}, while some of them (GRB 090510, 090902B, 090926A) clearly have an extra hard component at $\\gtrsim 10$~MeV~\\cite{Abd+09b,Ack+10,Ack+11}.  GRBs with the extra GeV component belong to the brightest class of GRBs, which may indicate that this could be more common, even though it can be seen only in bright LAT GRBs~\\cite{Gra10}. Another feature found by \\textit{Fermi} is that $\\gtrsim 100$~MeV gamma rays are delayed by $\\sim 0.1-1$~s in the cosmological rest frame, behind the onset of the MeV emission.  The origin of the high-energy emission is now under active debate. Late-time high-energy gamma-ray emission from GRBs such as 080916C, 090510, and 090902B has been attributed to afterglow emission rather than the prompt emission~\\cite{KD10,Ghi+10,Raz11}.  But, the high-energy emission in the early phase usually has a strong variability and often correlates with the MeV emission, which indicates an internal origin~\\cite{He+11}.  Since the simple one-zone leptonic model often has difficulties, multi-zone models, not only the SSC model~\\cite{CGP10,DBD11} but also the external inverse-Compton (IC) model~\\cite{TWM11,Pee+11} and the synchrotron model~\\cite{Iok10}, have been invoked.  The one-zone hadronic models are also viable, and the high-energy emission can be explained by proton-induced cascade or proton synchrotron emission~\\cite{AGM09,RDF10}.  The hadronic models are of interest, since GRBs may be the sources of observed ultra-high-energy cosmic rays (UHECRs), whose origin has been a big mystery for about 50 years~\\cite{Wax95,Vie95,MINN08}, and their neutrino signals may be seen by ongoing neutrino observations~\\citep[a recent review is][]{Wax11}.  In this paper, we study a slow heating scenario in the presence of electromagnetic cascades initiated by injection at high energies.  To explain the MeV emission self-consistenly, we consider effects of acceleration (via the second-order Fermi mechanism) or heating of secondaries, which can be anticipated in the post-shock region of relativistic shocks (or magnetic reconnection).  We show that cascades can play an important role in the injection of re-accelerated/heated particles, and slow heating scenarios via stochastic acceleration or wave heating typically lead to hard spectra compatible with observations.  As an injection process at high energies, we consider an application to hadronic injection via $pp$ and $p \\gamma$ reactions, whereby the hadronic models have the appealing feature of explaining an extra hard component at GeV energies.  The paper is organized as follows.  Section~2 describes an overview of the model, and we provide the numerical method and its results in Section~3.  We discuss implications for neutrinos in Sections~4, and our results are summarized in Section~6.  Throughout this work, cosmological parameters are set to $H_0=71~{\\rm km}~{\\rm s}^{-1}~{\\rm Mpc}^{-1}$, $\\Omega_M=0.3$, and $\\Omega_\\Lambda=0.7$, and we adopt the conventional notation $Q=Q_x \\times {10}^{x}$. ",
        "conclusions": "In this paper we studied the role of stochastic acceleration (or wave heating) on the GRB prompt emission, in the presence of electromagnetic cascades.  As an example, we employ the hadronic model to initiate cascade processes.  We demonstrated that some of the current issues in the GRB prompt emission can be explained.  (1) The low-energy photon index is expected to be $\\alpha \\sim 1$ when electrons and positrons confined in the downstream region are accelerated stochastically or heated via turbulence, and the typical peak energy, $E^{\\rm br}$, is stabilized by the balance between the stochastic acceleration (or heating) time and the cooling time. (2) It is possible to avoid the fast cooling problem because the charged leptons are slowly heated by turbulence.  When the leptons are injected via the hadronic processes, we may not have a potential issue on injection of primary electrons in the shock acceleration mechanism. (3) An extra hard component, $E F(E) \\propto E^{0-0.5}$, which has often been observed above $1-100$~MeV, can be explained by cascade emission, such as in proton-induced cascades.  The existence of the extra component is an important feature of the cascade initiated by high-energy injection well above MeV energies.  Although we have considered one-zone cases, dissipation will occur at various radii, which can affect the high-energy spectrum.  As shown in appendix B, in slow heating scenarios via stochastic acceleration, the high-energy photon spectrum can be harder than in the one-zone case because of the superposition of one-zone slow-heating spectra.  As a result, the broadband spectral shape can be a broken power law up to $\\sim$~GeV energies.  Such multi-zone effects are even more important above $\\gtrsim 0.1$~GeV.  Smaller $\\gamma \\gamma$ opacities at larger radii generally smear out the $\\gamma \\gamma$ pair-creation cutoff~\\cite{Aoi+10}.  Note that, in our calculation method (where spectra are obtained via iteration), the radiative transfer effect in the emitting slab, which leads to $(1-\\exp(-\\tau_{\\gamma \\gamma}))/\\tau_{\\gamma \\gamma}$ rather than the exponential cutoff, is not included.  Also, time-dependent and geometrical effects introduce additional complications, and the pair-creation break can be higher~\\cite{GTS08}.  Hence, we expect that the extra component extends to higher energies, although details are beyond the scope of this work.  One of the observational properties found by \\textit{Fermi} is the delayed onset at GeV energies.  It is not so obvious how to explain this in our model.  However, there are several reasons that may lead to broader pulses at high energies.  Especially, multi-zone effects seem important.  As expected from the above, the cascade component as well as the slow heated component would be superposed by dissipation at various radii.  Also, gamma rays from outer radii would be more transparent, which could make a dominant contribution at GeV-TeV energies.  In addition, in the hadronic models, neutron beams from the inner radius make further gamma rays via neutron decay and photomeson production with low-energy photons produced by dissipation at the outer radius.  Leptons may also up-scatter the MeV emission from inner radii via the external IC process. In addition, in our model, emissions on timescales longer than $t_{\\rm tur}$ and $t_{\\rm dyn}$ are expected.  Even after the dynamical time, the residual energy of electrons would be released by the external IC process and synchrotron radiation in some residual magnetic field.  In the hadronic models, further hadronic reactions may also happen as the escaping cosmic rays diffuse in the residual magnetic field (although they are adiabatically cooled).  Hence, future refinements based on multi-zone calculations seem necessary.  All hadronic models have typically several caveats or open issues.  They often lead to the requirement of a high baryon loading, which may be problematic, although this is motivated mainly by the GRB-UHECR hypothesis.  Another potential issue is the radiative efficiency.  In our cases, we typically obtain the reasonable radiative efficiency, $L_\\gamma/L_{\\rm th} \\sim 0.1$, although the apparent radiative efficiency can be higher due to lower $\\gamma_{\\rm reacc}$, release of the cosmic rays, residual emission and subsequent internal collisions.  Despite these concerns, there are interesting features.  Importantly, the hadronic model seems testable by upcoming multi-messenger observations.  The high-energy neutrino signal is one of the most important signatures, since the neutrino luminosity is expected to be comparable to the extra hard component.  If this extra component is ubiquitous and it carries $\\sim 1-10$~\\% of the total gamma-ray energy, IceCube would detect this signal in multi-year observations.  It is also important to check whether the extra hard component is common among all GRBs. As of now, its detectability with \\textit{Fermi} is limited and the current analyses suggest that the GeV emission has less than $\\sim 10$~\\% of the MeV emission.  Therefore, in order to unravel the properties of the high-energy emission, deeper observations via detections with Cherenkov detectors such as MAGIC and VERITAS should be important, even though detections from distant GRBs become difficult due to the attenuation by the extragalactic background light.  Although very-high-energy photons from GRBs have not been firmly detected so far, the future CTA~\\cite{CTA10} and HAWC are anticipated to change this situation.  We note that, although we have focused on the hadronic model as an example of high-energy injection, our results can be applied to other models, e.g., such as magnetic  dissipation.  In the slow heating scenario resulting from stochastic acceleration (or heating), it seems common to have hard low-energy photon indices, $\\alpha \\sim 1$ and $\\beta >2$. This feature itself would not be changed by details of the high-energy injection, as long as slow heating can operate and the MeV peak comes from balance between acceleration/heating and cooling. Hence, as for the origin of stochastically accelerated (or heated) particles, one may think of other high-energy injection processes and subsequent cascades.  If magnetic dissipation such as reconnection accelerates electrons up to very high energies and subsequent cascades are developed, similar spectra would be obtained."
    },
    "1107/1107.2163_arXiv.txt": {
        "abstract": "We consider a minimal scale-invariant extension of the Standard Model of particle physics combined with Unimodular Gravity formulated in \\cite{Shaposhnikov:2008xb}. This theory is able to describe not only an inflationary stage, related to the Standard Model Higgs field, but also a late period of Dark Energy domination, associated with an almost massless dilaton. A number of parameters can be fixed by inflationary physics, allowing to make specific predictions for any subsequent period. In particular, we derive a relation between the tilt of the primordial spectrum of scalar fluctuations, $n_s$, and the present value of the equation of state parameter of dark energy, $\\omega_{DE}^0$. We find bounds for the scalar tilt, $n_s<0.97$, the associated running, $-0.0006<d\\ln n_s/d\\ln k\\lesssim-0.00015$, and for the scalar-to-tensor ratio, $0.0009\\lesssim r<0.0033$, which will be critically tested by the results of the Planck mission. For the equation of state of dark energy, the model predicts $\\omega_{DE}^0>-1$. The relation between $n_s$ and $\\omega_{DE}^0$ allows us to use the current observational bounds on $n_s$ to further constrain the dark energy equation of state to $0< 1+\\omega_{DE}^0< 0.02$, which is to be confronted with future dark energy surveys. ",
        "introduction": "At the classical level, the Lagrangian describing the Standard Model of particle physics (SM) minimally coupled to General Relativity (GR) contains three dimensional parameters: Newton's constant G, the vacuum expectation value (vev) of the Higgs field or, equivalently, the Higgs boson mass and a possible cosmological constant $\\L$. The masses of quarks, leptons and intermediate vector bosons are induced by the vev of the Higgs field. At the quantum level, additional scales, such as $\\Lambda_\\textrm{QCD}$ and all other scales related to the running of coupling constants, appear due to dimensional transmutation. It is tempting to look for models in which all these seemingly unrelated scales have a common origin.  In this work we present a detailed analysis of a model realizing this idea, proposed in \\cite{Shaposhnikov:2008xb}. We will refer to it as the Higgs-Dilaton model. The model is based on a minimal extension of the SM and GR that contains no dimensional parameters in the action and is  therefore scale-invariant at the classical level. Scale invariance is achieved by introducing a new scalar degree of freedom, called Dilaton. The motivation of the model relies on the assumption that the structure of the theory is not changed at the quantum level. In other words, the full quantum effective action should still be scale-invariant and the effective scalar potential should preserve the features of the classical potential. A perturbative quantization procedure maintaining scale invariance was presented in \\cite{Shaposhnikov:2008xi} (see also \\cite{Englert:1976ep}). In the Higgs-Dilaton model all scales are induced by the spontaneous breakdown of scale invariance (SI). As a consequence of the broken symmetry, the physical dilaton is exactly massless. Replacing GR by Unimodular Gravity (UG), in which the metric determinant is fixed to one, ${|g|=1}$, results in the appearence of an arbitrary integration constant in the equations of motion, representing an additional breaking of scale symmetry. As discussed in \\cite{Shaposhnikov:2008xb}, in theories with scalar fields non-minimally coupled to gravity, this constant effectively gives rise to a non-trivial potential for the scalar fields. In the case of the Higgs-Dilaton model, the new potential is of the ``run-away'' type in the direction of the dilaton.  While the dynamical breakdown of the scale symmetry by the Higgs field can provide a mechanism for inflation in the early universe \\cite{Bezrukov:2007ep}, the light dilaton, practically decoupled from all SM fields, can act as Quintessence (QE), i.e. as dynamical Dark Energy (DE). We find that, under some assumptions, it is possible to relate the observables associated to inflation to those associated to dark energy. Namely, we establish a functional relation between the predicted value for the tilt $n_s$ of the primordial scalar power spectrum and the predicted equation of state parameter $w_{DE}^0$ of dark energy. Further, we find a relation involving the corresponding second order quantities, i.e the running $\\alpha_\\zeta$ of the spectral tilt and the rate of change $w_{DE}^a$ of the DE equation of state.  The paper is organized as follows: In section \\ref{msi} we introduce and discuss the minimal scale-invariant extension of the Standard Model and General Relativity. In Section \\ref{siug}  the idea of Unimodular Gravity is described and applied to the scale-invariant model. We then discuss the cosmology of the resulting Higgs-Dilaton Model. The inflationary period is studied in detail in Section \\ref{hdi}. The implications of the model for the late dark energy dominated stage are studied in Section \\ref{lu}. Finally, conclusions are presented in Section \\ref{conclusions}. For completness, an analysis of slow-roll inflation in the Jordan frame and the differences with respect to the Einstein frame are presented in Appendix \\ref{appendix}. ",
        "conclusions": "We have considered a minimal scale-invariant extension of the Standard Model, non-minimally coupled to gravity, including a scalar dilaton. All mass scales at the classical level, including the Planck scale and the Electroweak scale, are induced by the spontaneous breaking of the scale invariance. The physical dilaton is almost massless but hardly affects particle physics phenomenology. Our findings rely on SI both at the classical and the quantum level \\cite{Shaposhnikov:2008xi,Englert:1976ep}. The replacement of standard General Relativity by Unimodular Gravity gives rise to an arbitrary constant in the equations of motion, which in a minimally coupled theory would play the role of a cosmological constant. However, due to the non-minimal couplings between the scalar and the gravitational sectors, this constant gives rise to a non-trivial ``run-away'' potential for the dilaton. As a consequence, the dilaton can play the role of a quintessence field, responsible for a late dark energy dominated stage. For appropriate values of the free parameters and initial conditions, the constructed model presents a rich cosmological phenomenology, providing mechanisms both for inflation and dark energy.  We find that the amplitude $P_\\zeta$ of CMB anisotropies depends mainly on the ratio $\\xi_h/\\sqrt\\lambda$ while the spectral tilt $n_s$, the associated running $\\alpha_\\zeta$ as well as the scalar-to-tensor ratio $r$ depend mainly on $\\xi_\\chi$. The observational limits on $P_\\zeta$ and $n_s$ put bounds on $\\xi_h/\\sqrt\\lambda$ and $\\xi_\\chi$, which in turn provide the bounds $\\alpha_\\zeta\\lesssim-0.00015$ and $r\\gsim 0.0009$. In addition, the model predicts the bounds $n_s<0.97$, $\\alpha_\\zeta>-0.0006$ and $r<0.0033$, which are obtained in the limit $\\xi_\\chi\\rightarrow 0$ and correspond to the predictions of the Higgs-Inflation model of \\cite{Shaposhnikov:2008xb}. The confrontation of these bounds with the results of the Plack satellite mission will constitute an important test of the Higgs-Dilaton model.  Neither SI nor Unimodular Gravity forbid the existence of a quartic term $\\beta\\chi^4$ in the Jordan frame, which would correspond to a proper cosmological constant in the Einstein frame. However, the parameter choice forbidding such a term $(\\beta=0)$ appears to be specially interesting, both from the cosmological and the quantum theory point of view. For this choice, the dilaton is alone responsible for dark energy. The associated equation of state parameter $w_{DE}^0$ is found to practically depend on $\\xi_\\chi$ only. This has the interesting consequence that the spectral index $n_s$ can be expressed as $n_s=n_s(w_{DE}^0)$ thus relating an observable from the very early universe to an observable of the present universe. For a particular parameter region, this relation takes the simple form $-3(w^0_{DE}+1)\\simeq(n_s-1)$. The observational bound on $n_s$ translates into a bound $0\\leq1+w^0_{DE}\\lesssim 0.02$, which might be tested by future experiments. Further, we were able to derive a relation between the running of the spectral index $\\alpha_\\zeta$ and the rate of change $3w^a_{DE}$ of the equation of state parameter. Notice that for the dilaton to provide the measured abundance of dark energy, initial conditions have to be finely tuned. Hence, as is the case for all quintessence models, the Cosmological Coincidence Problem remains unsolved."
    },
    "1107/1107.5427_arXiv.txt": {
        "abstract": "Several recent studies have shown how to properly calculate the observed clustering of galaxies in a relativistic context, and uncovered corrections to the Newtonian calculation that become significant on scales near the horizon. Here, we retrace these calculations and show that, on scales approaching the horizon, the observed galaxy power spectrum depends strongly on which gauge is assumed to relate the intrinsic fluctuations in galaxy density to matter perturbations through a linear bias relation. Starting from simple physical assumptions, we derive a gauge-invariant expression relating galaxy density perturbations to matter density perturbations on large scales, and show that it reduces to a linear bias relation in synchronous-comoving gauge, corroborating an assumption made in several recent papers.  We evaluate the resulting observed galaxy power spectrum, and show that it leads to corrections similar to an effective non-Gaussian bias corresponding to a local $f_{\\rm NL,eff} \\lesssim 0.5$. This number can serve as a guideline as to which surveys need to take into account relativistic effects.  We also discuss the scale-dependent bias induced by primordial non-Gaussianity in the relativistic context, which again is simplest in synchronous-comoving gauge. ",
        "introduction": "\\label{sec:intro}  The clustering of galaxies and other large-scale structure (LSS) tracers on the largest scales has recently received great interest as a probe of inflation and its alternatives.  In the presence of primordial non-Gaussianity, biased tracers can exhibit a significant scale-dependent bias with respect to the matter distribution which increases strongly towards large scales \\cite{DalalEtal08}.  This can be used as a sensitive probe of primordial non-Gaussianity \\cite{SlosarEtal}.  Furthermore, ongoing, future, and proposed surveys such as the Baryon Oscillation Spectroscopic Survey \\cite{BOSS}, the Hobby Eberly Telescope Dark Energy Experiment \\cite{hetdex}, HyperSuprime Cam, the Dark Energy Survey\\footnote{\\tt http://www.darkenergysurvey.org/}, the Subaru Prime Focus Spectrograph, BigBOSS \\cite{BigBOSS}, the Large Synoptic Survey Telescope\\footnote{\\tt http://www.lsst.org/lsst}, the Wide-Field Infrared Survey Telescope\\footnote{\\tt http://wfirst.gsfc.nasa.gov/}, Euclid\\footnote{\\tt http://sci.esa.int/euclid/}, and the Square Kilometer Array\\footnote{\\tt http://www.skatelescope.org/} will probe modes that approach the comoving horizon.  All of this is strong motivation to go beyond the Newtonian picture of galaxy clustering widely adopted so far, and to embed this observable into a proper relativistic context.  This is analogous to what has been done for the cosmic microwave background (CMB), and several aspects have long been worked out \\cite{SW67,Sasaki87,BonvinEtal06}.  However, galaxy clustering involves a few additional complications: first, it is intrinsically three- rather than two-dimensional; second, one has to take into account selection effects such as cuts on observed flux and redshift; and third, the relation between intrinsic fluctuations in the galaxy density and the fluctuations in the matter density is non-trivial.  In order to use the clustering of LSS tracers on scales approaching the horizon $c/H(z)$, we thus need to understand the connection between the theoretical predictions for cosmological perturbations and the observationally inferred overdensities of galaxies.  In particular, the perturbations in the metric, matter density, velocity, etc. are always defined with respect to a particular choice of coordinates (gauge), whereas observables should be independent of this gauge choice.  Recently, \\citet{yoo/etal:2009} have shown how to calculate the observed galaxy overdensity in a generally covariant, relativistic context.  In this paper, we perform a similar derivation, generalizing their results in one important aspect.  On sufficiently large scales so that perturbations are linear (and assuming Gaussian initial conditions), one commonly assumes a linear relation between galaxy overdensities and perturbations in the matter density.  As we show here however, the observed galaxy power spectrum depends on which gauge these overdensities are referred to.  For example, the linear bias assumed in \\cite{yoo/etal:2009,yoo:2010} is equivalent to a linear bias relation in the constant-observed-redshift gauge;  on the other hand, \\cite{challinor/lewis:2011,BruniEtal} adopted a linear bias in synchronous-comoving gauge.  Clearly, this situation is not satisfactory, since the gauge choice should not impact \\emph{any observable} quantity.  However, we can make progress using simple physical arguments.  In a universe with Gaussian adiabatic perturbations, a galaxy knows about two properties of its large-scale environment: the average, ``background'' density of matter, and the local age (growth history) of its environment. Thus, a general bias expansion should involve \\emph{both} density and age (or local growth factor).  In this context, the simplest gauge choice is the synchronous-comoving (sc) gauge, where constant-time hypersurfaces are equivalent to constant-age hypersurfaces.  Then, the bias with respect to age becomes irrelevant, and we recover the well-known linear bias relation: $\\d^{(\\rm sc)}_g = b\\:\\d^{(\\rm sc)}_m$.  Further advantages of synchronous-comoving gauge are that the density field in N-body as well as the output of commonly used Boltzmann codes are given in this gauge.  We shall thus express most of our results in synchronous-comoving gauge.  The transition to other gauges can be performed easily using expressions given in the appendices.  Note that when properly transformed, the results derived in different gauges should agree.  Recent papers by \\citet{challinor/lewis:2011} and \\citet{bonvin/durrer:2011} provide a further reason to reinvestigate relativistic corrections to the observed galaxy correlation, as they were not able to reach agreement with the expressions given in Ref.~\\cite{yoo/etal:2009}.  The outline of the paper is as follows.  We begin by deriving the observed galaxy density in terms of perturbations in the synchronous-comoving gauge in \\refsec{ng}.  In \\refsec{bias}, we discuss how galaxy biasing can be implemented in a gauge-invariant way. \\refsec{Pkg} discusses the observed galaxy power spectrum. We conclude in \\refsec{concl}.  In the appendices, we present useful results on the conversion between different gauges and metric conventions (\\refapp{gTransf}), more details on the derivations (\\refapp{Pgderiv}), and various analytical test cases for the expression for the observed galaxy overdensity (\\refapp{test}).  We also make the connection with the work of other recent papers \\cite{yoo/etal:2009,yoo:2010,challinor/lewis:2011,bonvin/durrer:2011} in \\refapp{comp}. ",
        "conclusions": "\\label{sec:concl}  Future galaxy surveys will measure the clustering of galaxies on scales approaching the horizon.  Thus, it is necessary to embed the observed galaxy density in a relativistic context.  This problem has received considerable attention recently.  In this paper, we have derived the observed galaxy density contrast $\\tilde\\d_g$ in terms of the intrinsic galaxy overdensity $\\d_g$ and metric perturbations in the synchronous-comoving gauge (\\refsec{ng}). By transforming to conformal-Newtonian gauge, we reach agreement with the results of Refs.~\\cite{challinor/lewis:2011,bonvin/durrer:2011}. On the other hand, we find some minor disagreement with the expression of \\cite{yoo/etal:2009}, which can be traced back to a sign issue. In \\refapp{test}, we also show that our formula for $\\tilde\\d_g$ reproduces the correct analytic result for the following five test cases: 1) pure spatial gauge mode, 2) zero-wavenumber gauge mode, 3) perturbation to the expansion history, 4) spatial curvature, 5) Bianchi I cosmology.  These tests cover most of the parameter space of actual cosmologies in a simplified setting, and thus lend credibility to our result.  One necessary, further ingredient for a description of galaxy clustering on large scales is a physical, gauge-invariant definition of galaxy bias.  We present a straightforward physically motivated definition in \\refsec{bias}.  This bias relation is easily seen to reduce to the standard linear bias relation in synchronous coordinates, where constant-time hypersurfaces coincide with constant-age hypersurfaces. The bias relation \\refeq{bias1} can be seen as a proper generalization of the peak-background split bias.  Using this result, we arrive at a simple expression for the observed galaxy density perturbations in synchronous-comoving gauge [\\refeq{dgreal}], described by the (gauge-invariant) bias parameter $b$, and the redshift evolution of the tracer population [through $b_e$, \\refeq{btaudef}].  The recent study of \\citet{BaldaufEtal} has also reached the same conclusion by constructing locally flat space-time coordinates around a freely-falling observer, where long-wavelength modes locally act as curvature. In that coordinate system, the galaxy number density is modulated not only by the local curvature but also the time difference between global time and local time whose effect exactly coincides with our $\\bt$ in \\refeq{bias1}. In contrast, Ref.~\\cite{yoo/etal:2009} adopted a bias relation in the constant-observed-redshift gauge, which leads to considerably different predictions for the large-scale galaxy power spectrum. However, this does not seem to be a physical description of galaxy bias, since the age of the Universe is \\emph{not} constant on constant-observed-redshift slices.  We believe these results, together with other recent work \\cite{challinor/lewis:2011,bonvin/durrer:2011,BaldaufEtal}, allow us to unambiguously predict the two-point statistics of large-scale structure tracers on large scales."
    },
    "1107/1107.0728_arXiv.txt": {
        "abstract": "Accretion disks that become gravitationally unstable can fragment into stellar or sub-stellar companions. The formation and survival of these fragments depends on the precarious balance between self-gravity, internal pressure, tidal shearing, and rotation. Disk fragmentation depends on two key factors (1) whether the disk can get to the fragmentation boundary of Q=1, and (2) whether fragments can survive for many orbital periods.  Previous work suggests that to reach Q=1, and have fragments survive, a disk must cool on an orbital timescale. Here we show that disks heated primarily by external irradiation always satisfy the standard cooling time criterion. Thus even though irradiation heats disks, and makes them more stable in general, once they reach the fragmentation boundary, they fragment more easily. We derive a new cooling criterion that determines fragment survival, and calculate a pressure modified Hill radius, which sets the maximum size of pressure-supported objects in a Keplerian disk. We conclude that fragmentation in protostellar disks might occur at slightly smaller radii than previously thought, and recommend tests for future simulations that will better predict the outcome of fragmentation in real disks. ",
        "introduction": "The past two decades have witnessed growing interest in the conditions under which gravitationally unstable disks fragment to produce bound objects. These fragments may grow into massive stars or black holes within AGN disks, or small stellar and substellar companions in protostellar disks. In either context, we can divide the process into two steps. First, the disk must become linearly gravitationally unstable, and secondly, once formed, the fragments must be able to survive the various disruptive forces in the disk.  The first step, linear gravitational instability (GI),  requires that the disk be Toomre unstable: $ Q= c_s \\Omega / \\pi G \\Sigma \\leq Q_0 \\approx 1$ \\citep{Toom1964,Saf1960},  where $c_s = \\sqrt{kT/\\mu}$ is the isothermal sound speed of the gas with mean particle weight $\\mu$ and temperature $T$, $G$ is the gravitational constant, $\\Sigma$ is the disk surface density, $k$ is the Boltzmann constant,  and $\\Omega$ is the orbital angular frequency.  The second stage is more uncertain. Three primary forces compete to destroy newly born fragments: internal gas pressure, tidal forces from the central object  (shearing), and internal rotation. Fragments can also destroy each other through mutual interactions.  Whether the disk can be driven to the fragmentation boundary, which we call $Q=Q_0$, and whether fragments can survive, must be considered separately.  Both requirements place limits on the cooling properties of disks. Moreover, the answer to both of these questions depends on whether the disk is actively heated (i.e. by dissipation of accretion energy through an effective turbulent $\\alpha$ viscosity), or passively heated (i.e. be external irradiation).  Both processes operate in most disks, but it matters which is the dominant energy source. Here we consider the passive (irradiated) case in detail, as it has been neglected in previous theoretical studies. { Throughout the text we use the term viscosity to describe any process that contributes to both transport of angular momentum and local energy dissipation.}  The role of cooling in fragmentation has been discussed in the literature in two ways, which we dub the Viscous Criterion \\citep{Pringle81,Gam2001}  and the Collisional Criterion \\citep{1989ApJ...341..685S}. The Viscous Criterion describes whether an actively-heated disk can rid itself of the energy generated by GI-driven accretion on the disk's dynamical timescale, while the Collisional Criterion considers whether newborn fragments will collide with each other on this timescale. While the Viscous Criterion addresses the first question: (1)  Can a disk be driven to $Q=Q_0$ so that initial collapse begins, the Collisional Criterion addresses the second: (2) Can a collapsing fragment survive?  Here we provide a more complete picture of the necessary disk conditions for fragmentation and fragment survival.  We begin with question (1), emphasizing the crucial role of disk history in determining whether a disk can be driven to collapse.  The Viscous Criterion answers this question solely for viscously-heated disks, and we provide a complementary criterion for irradiated disks, which we show automatically satisfy the Viscous Criterion.  We then turn to question (2).  We find that the Collisional Criterion does not fully capture the role of cooling in fragment survival, because it does not account for the contraction of fragments.  We propose a more general cooling criterion which governs fragment survival in all disks. This ``Stalling Criterion\" first determines whether fragments collapse on the free-fall timescale, or if they quickly become pressure supported.  If a fragment becomes pressure supported, stalling collapse, it must be smaller than a pressure-modified Hill radius to avoid tidal disruption.  Fragments that contract to small enough sizes before stalling can survive while slowly undergoing Kelvin-Helmholz contraction.  Fragment survival depends both on the cooling time and the gas equation of state. The Stalling Criterion is applicable in both viscously-heated and irradiated disks. \\Fig{fig-3crit} illustrates how the different cooling criteria we describe determine whether disks fragment, and whether their fragments survive.  We begin with brief definitions of the disk cooling time and related terms in Section \\ref{coolingtimes}. We then distinguish quantitatively between disks heated via external irradiation or internal viscous dissipation in Section \\ref{sec-irradmean} and describe the run-up to instability in each case (Sections \\ref{sec-viscdrive} and \\ref{sec-irraddrive}). We introduce our new cooling criterion in Section \\ref{sec-KHC}. We show that for arbitrary cooling time, the Stalling Criterion sets a critical size scale ($R_{\\rm H+p}$) below which a fragment can survive if pressure supported.  Whether or not a fragment can reach this critical radius depends on the disk equation of state.  We review other means for disruption of fragments in Section \\ref{caveats}.  We discuss the implications of these processes for protostellar and AGN disks in Section \\ref{sec-irrad}, then comment briefly on recent simulations of viscously-heated disks (Section \\ref{sec-visc}).   Finally, we provide suggestions for future simulations of irradiated disks and conclude in Section \\ref{discussion}. ",
        "conclusions": "\\label{discussion}  In this paper, we separate the conditions for successful fragmentation into two considerations: (1) Can the disk reach the threshold of linear instability? and (2) Can a fragment survive after it has begun to collapse? This is distinct from previous work that has focused on whether or not the disk can reach $Q=1$ {\\emph{ and}} cool quickly. We show that cooling is intimately connected both to reaching instability and to fragment survival.  We first address the distinction between disks heated internally by the dissipation of accretion energy, and those heated externally by irradiation. We show that the Viscous Criterion only governs the answer to question (1) in viscously heated disks because irradiated disks always effectively satisfy the Viscous Criterion. We show that fragmentation is easier in irradiated disks, and that if an irradiated disk is thin enough (or equivalently low enough in mass) when $Q\\sim1$, global transport may be ineffective, allowing the disk to be driven to collapse by low rates of mass infall. Thus, fragmentation may occur at slightly smaller radii, and lower initial masses then previously assumed.  Importantly, fragmentation at low accretion rates could reduce the expected growth of fragments after their formation.  We then derive a generic cooling criterion which controls fragment survival in all disks. The Stalling Criterion divides fragments that can collapse on their own free-fall timescale from those that will become pressure supported and begin Kelvin-Helmholtz contracting. Fragments in the latter category must contract to a critical radius ($R_{\\rm H+p}$) to avoid being torn apart by tides when pressure stalls their collapse. This radius is roughly half the Hill radius: the reduction is due to the inclusion of pressure support in the calculation of force balance.  For reasonable perturbations, fragments in disks with soft equations of state ($\\gamma =7/5$) will likely be able to contract below $R_{\\rm H+p}$, {\\emph{even}} if they cool over many orbital periods.  In contrast, for stiffer equations of state, fragments that do not collapse in free fall will likely be sheared apart.  We note that the value of $R_{\\rm H+p}$ depends on fragment geometry, and whether a given fragment will contract to this size depends on its non-linear evolution in time. For example, if fragmentation occurs at densities and temperatures where there are sharp jumps in opacity, this must be taken into account in determing $\\beta$.  Specific contraction models for disk-born fragments at all masses should be compared to the Stalling Criterion. We comment on other disruptive processes, including the classic Collisional Criterion, and the effects of fragment rotation, in Section \\ref{caveats}.  Of course, since fragments will continue to grow following formation, gravitational instability should still favor the production of low-mass brown dwarfs over planets \\citep{KMCY10}; That GI produces larger objects is increasingly supported by observations of M-star and brown-dwarf companions to A-stars \\citep{2010ApJ...712..421H}, and the paucity of wide, planetary mass companions \\citep{Leconte:2010}.  \\subsection{Suggestions for Future Work} This work should be a useful guide for future numerical simulations intended to study fragmentation in irradiation-dominated disks.  Two major uncertainties limit our conclusions for protostellar disks: \\begin{enumerate} \\item What is the value of $\\alpha_{\\rm sat}$ in a viscously-heated $Q\\sim 1$ disk? \\item What is the maximum GI-driven transport rate as a function of $H/r$ in an irradiated $Q\\sim1$ disk? \\end{enumerate} These two uncertainties are effectively the same question, but since the physics of transport in these regions differs, they must be considered independently.  Simulations will also be required to confirm our conclusion that regardless of whether they collapse in free fall, fragments can likely survive if $\\gamma = 7/5$ because they collapse to smaller than $R_{\\rm H+p}$ on of order a free-fall time. Hence, disks with soft equations of state must simply reach $Q=Q_0$ to fragment.  Confirmation of this conclusion will require a full treatment of radiative transfer and disk opacity.  As noted earlier, the impact of fragment rotation must also be explored. { Current state of the art simulations of disks with radiative transfer do not yet cover the range of parameter space required to answer the above questions (\\citealt{Boss:2007,2008ApJ...673.1138C, Cai:2010,Boley2010}; note that the latter two authors have questioned the conclusions of the former based on differing treatments of the photosphere).}  Answers from such simulations will enable us to distinguish between the following outcomes in each regime. \\begin{itemize} \\item {\\it Viscously-heated regions:} If $\\alpha_{\\rm sat} \\gtrsim 0.1$, then standard applications of the cooling time criterion give accurate results.  If $\\alpha_{\\rm sat} \\lesssim 0.1$, then the standard Viscous Criterion will determine the boundary for fragmentation, but the critical cooling time will be longer. \\item {\\it Irradiation-dominated regions:} If  both the maximum GI-driven transport rate and $\\alpha_{\\rm eff}$ due to MRI  are too low to process the infall rate onto a protostellar disk, fragmentation may be possible for thin disks at infall rates lower than $\\sim$$c_s^3/G$.  If this occurs, low-mass fragments which do not experience substantial growth might be produced.  Otherwise, high accretion rates are required to drive these disks unstable. \\end{itemize}"
    },
    "1107/1107.1518_arXiv.txt": {
        "abstract": "We perform two-dimensional, Point-Spread-Function-convolved, bulge+disk decompositions in the $g$ and $r$ bandpasses on a sample of 1,123,718  galaxies from the Legacy area of the Sloan Digital Sky Survey Data Release Seven. Four different decomposition procedures are investigated which make improvements to sky background determinations and object deblending over the standard SDSS procedures that lead to more robust structural parameters and integrated galaxy magnitudes and colors, especially in crowded environments. We use a set of science-based quality assurance metrics namely the disk luminosity-size relation, the galaxy color-magnitude diagram and the galaxy  central (fiber) colors  to show the robustness of our  structural parameters. The best procedure utilizes simultaneous, two-bandpass decompositions. Bulge and disk photometric errors remain below 0.1 mag down to bulge and disk magnitudes of  $g \\simeq 19$ and $r \\simeq 18.5$. We also use and compare three different galaxy fitting models: a pure S\\'ersic model, a $n_b=4$ bulge + disk model and a S\\'ersic (free $n_b$) bulge + disk model. The most appropriate model for a given galaxy is determined by the $F$-test probability. All three catalogs of measured structural parameters, rest-frame magnitudes and colors are publicly released here. These catalogs should provide an extensive comparison set for a wide range of observational and theoretical studies of galaxies. ",
        "introduction": "The properties of disks and spheroids in the local Universe are the direct outcomes of the hierarchical assembly of galaxies over cosmological time. It is well known that the structural parameters of these important galaxy sub-components obey scaling relations between size, luminosity and internal velocity \\citep[e.g.,][]{courteau07}. The amount of stellar mass found in disks and bulges places strong constraints on the galaxy merger tree from $\\Lambda$CDM N-body simulations \\citep{hopkins10} and the secular formation of bulges \\citep[e.g.,][]{kormendy79,athanassoula03}. Some properties such as disk size can be traced to properties of host dark matter haloes such as specific angular momentum \\citep{fall80,mmw98}, and more information is needed for a truly large sample of galaxies to improve the recipes used in semi-analytical approaches to the treatment of the so-called ``sub-grid\" physics \\citep{governato04,robertson04,dutton09}.  Thanks to the availability of large datasets and powerful computing clusters, it is now possible to analyze large samples of galaxies in a fully quantitative way. Quantitative measurements of galaxy properties have made significant progress over the last few years \\citep{peng02,simard02,desouza04,lotz04,conselice06,pignatelli06} and now provide an important framework for comparisons between observation and theory. A number of previous works have looked at the morphologies of SDSS galaxies using a wide range of subsamples and techniques \\citep{shen03,blanton05a,kelly05,benson07,fukugita07,lintott08,labarbera10,wijesinghe10}. This paper presents the largest catalog of bulge+disk structural parameters measured to date of galaxies in the Sloan Digital Sky Survey Data Release Seven \\citep[SDSS DR7,][]{abazajian09}. This catalog has already been used to study visual versus quantitative morphologies \\citep{cheng10}, galaxy pairs \\citep{ellison08,ellison10, patton11}, disk scaling relations at low and high redshift \\citep{dutton11a,dutton11b}, the disk size function \\citep[][and Simard 2011, in preparation]{kanwar08}, the evolution of cluster galaxies with redshift \\citep{simard09}, compact groups of galaxies \\citep{mendel11}, and the disk and bulge stellar mass functions \\citep{thanjavur11}. The paper is organized as follows. The data and the bulge+disk decompositions are described in Sections~\\ref{data} and~\\ref{analysis}.  The catalog, science-based quality assessment metrics and a comparison between different fitting models are presented in Section~\\ref{results}. The cosmology adopted throughout this paper is ($H_0, \\Omega_{m}, \\Omega_{\\Lambda}$) = (70 km s$^{-1}$ Mpc$^{-1}$, 0.3, 0.7). ",
        "conclusions": "\\label{summary} We have performed bulge+disk decompositions for 1.12 million galaxies in the Legacy area of the SDSS Data Release 7. Four decomposition procedures were used, and one of these procedures clearly produced more robust structural parameters, magnitude and colors when looking at three science-based data quality assurance metrics. The most reliable procedure included the following three important steps: \\begin{itemize} \\item GIM2D-based sky background level determination \\item SExtractor object deblending \\item Simultaneous bulge+disk decomposition in $g$ and $r$ \\end{itemize} We also used three different fitting models: a $n_b=4$ bulge + disk model, a free-$n_b$ bulge + disk model and a pure S\\'ersic model, and we provide a detailed comparison between the measured structural parameters from these fitting models with a mean to select the appropriate model for a given galaxy using the $F$-statistic. This comparison highlights the importance of this selection for a given science goal. Pure S\\'ersic model fits might be better suited to studies of global galaxy colours whereas a cut on the $F$-test probability $P_{pS}$ to select galaxies for which a bulge+disk model was required would be better for studies of bulge/disk colours. Using a low cut on both $P_{pS}$ and $P_{pS}$ probabilities would select relatively small but very robust samples of bulge+disk decompositions. Full catalogs of structural parameters are included here with important cautionary notes on how to select bulge and/or disk subsamples, how to select different morphological types, the issue of bars and the issue of internal dust. These catalogs should provide an extensive comparison set for a wide range of observational and theoretical studies of galaxies."
    },
    "1107/1107.4084_arXiv.txt": {
        "abstract": "We develop a general theory for estimating the probability that a galaxy cluster of a given shape exists. The theory is based on the observed result that the distribution of galaxies is very close to quasi-equilibrium, in both its linear and nonlinear regimes. This places constraints on the spatial configuration of a cluster of galaxies in quasi-equilibrium. In particular, we show that that a cluster of galaxies may be described as a collection of nearly virialized subclusters of approximately the same mass. Clusters that contain more than 10 subclusters are very likely to be completely virialized. Using our theory, we develop a method for comparing probabilities of different spatial configurations of subclusters. As an illustrative example, we show that a cluster of galaxies arranged in a line is more likely to occur than a cluster of galaxies arranged in a ring. ",
        "introduction": "Clusters and groups of galaxies are commonly studied structures in the universe, and have been defined using various criteria; groups are smaller clusters. These clusters may contain over a thousand member galaxies and occupy a few cubic megaparsecs. Although clusters can be identified by their density and size~(e.g. \\citealt{1958ApJS....3..211A, 1957PASP...69..409H}) their shapes and structures differ. For example, \\citet{1957PASP...69..409H} identify three classes of clusters: \\textit{compact clusters} which have a single nearly spherical dense concentration of galaxies, \\textit{medium compact clusters} which are less dense and may have multiple concentrations of galaxies and \\textit{loose clusters} which do not have any outstanding concentrations of galaxies.  Even these simple classes of clusters suggest greater complexity than just spherical concentrations in a region of space. For example, \\citet{1987AJ.....94..251B} found significant substructure in the core of the Virgo cluster as well as pronounced double structure, although \\citet{1957PASP...69..409H} classify the Virgo cluster as a medium compact cluster. The structure of the Virgo cluster, as the nearest large cluster, suggests that these irregular shapes are common.  Such irregular shapes may result from mergers. Smaller groups fall into the central region of a cluster and form subgroups whose member galaxies are still tightly bound to each other. Irregular shapes resulting from subgroups then disappear as a cluster virializes. However, many clusters have dynamical relaxation timescales on the order of a Hubble time, and their incomplete virialization suggests that irregular clusters with multiple concentrations of galaxies should be common in the universe.  The basic dynamical description of a cluster is its 6-dimensional phase space configuration such as a sphere with a density and velocity profile. More detailed descriptions of clustering include correlation functions, percolation trees and counts-in-cells statistics. In particular, the counts-in-cells description is especially suitable for this problem because it straightforwardly analyzes regions of space~(cells) with a specified size and shape. In addition, the physics of this description can be derived from gravitational thermodynamics~\\citep{1984ApJ...276...13S} or statistical mechanics~\\citep{2002ApJ...571..576A} where the basic particles of the system are galaxies that interact in a grand canonical ensemble of cells.  This physical description leads to the gravitational quasi-equilibrium distribution~(GQED) for the counts-in-cells distribution of galaxies. The GQED has been shown to agree with $N$-body simulations~(\\citealt{1988ApJ...331...45I} and subsequent work) and various analyses of sky surveys in both the linear and nonlinear regimes of clustering~(e.g. \\citealt{2005ApJ...626..795S,2011ApJ...729..123Y} and references therein). This indicates that the statistical mechanical theory is a suitable description of clustering.  The theory was subsequently extended to describe particles with different masses~\\citep{2006IJMPD..15.1267A}, three body interactions~\\citep{2006ApJ...645..940A} and different internal structures~\\citep{2010arXiv1011.0176Y}. It was also further generalized to take into account higher orders in the series expansion of the gravitational interaction~\\citep{2010ApJ...720.1246S}, and extended by \\citet{2004ApJ...608..636L} to describe the potential and kinetic energies in a single cell, and hence the probability that a cell with $N$ galaxies would be virialized.  In this paper, we extend the work by \\citet{2004ApJ...608..636L} and relate the ratio of potential energy to kinetic energy in a cell to its detailed configuration. Here, we work with cells rather than clusters because cells are well-defined regions of space and, unlike clusters, they have a clear boundary. Another advantage of working with cells is that we can use larger cells to study how clusters of galaxies cluster around each other, and hence provide insights into the process of hierarchical structure formation.  This paper is structured as follows: In section \\ref{sec-STP} we rescale the thermodynamic variables to describe the properties of a self-gravitating system in a form that is independent of physical units. This provides a scale-free description of particles in a cell. These may be galaxies in a group, or substructures in a cluster, or even clusters in a large cell for insights into hierarchical clustering. In section \\ref{sec-PE} we use the scaled thermodynamic variables to determine the probability of finding a cell with a given kinetic energy or virial ratio in quasi-equilibrium. This lets us determine the conditions that produce a virialized or unvirialized cell, and its implications for substructure in these cells. In section \\ref{sec-config} we derive a relation between the configuration, energy and probability of a cell. In section \\ref{sec-desc} we describe a procedure for studying the internal structure of a cell and apply this procedure to an illustrative example. Finally in section \\ref{sec-conc}, we summarize our results. ",
        "conclusions": "\\label{sec-conc}  In this paper we have described a theory of clustering that extends the work of \\citet{2004ApJ...608..636L} to calculate the probability that a cell has a given time-averaged virial ratio $\\overline{\\psi}$. This probability is independent of the average number of galaxies in a cell and depends on the clustering parameter $b$, the number of galaxies $N$ in the cell and the range of virial ratios that the cell may have.  A feature of this theory is the use of subclusters to describe the internal spatial configuration of a cell. These subclusters are concentrations of galaxies that are nearly virialized and share a common bulk velocity. The use of subclusters considerably simplifies the analysis because their masses are approximately equal and most clusters that are still assembling will have less than 10 subclusters. This is because cells with more than 10 subclusters are likely to already be virialized, and subclusters that are less massive than a tenth of the most massive subcluster do not have enough mass to significantly affect the potential energy.  Because subclusters are collections of galaxies, the number of subclusters within a single cluster can be an indicator of its evolutionary stage. In hierarchical models, clusters build up by the merger of multiple subclusters. Thus the presence of multiple clusters within a cell, or significant substructure within a cluster would suggest that mergers were relatively recent. Such clusters, having multiple subclusters, are classified as medium compact clusters in the classification scheme of \\citet{1957PASP...69..409H}. More evolved relaxed clusters have less substructure with a single dominant concentration. Such clusters are their own subcluster and are classified as compact clusters under the same classification scheme.  In addition, clusters with a negative specific heat may take on almost any configuration because $W_*[\\psi]$ is double-valued with a minimum at $\\psi=2/3$. Hence for almost every value of $W_*$ there is a value of $\\psi$ in the negative specific heat branch. The only spatial configurations that are unlikely to occur in the negative specific heat branch are those with very negative values of $W_*$. Since the absolute value of $W_*$ depends on the inverse separation between particles, such configurations will have particles that are very close to each other, and therefore are likely to be unstable to mergers. Together with the roughly $20\\%$ statistical variation over time of $\\psi$ about $\\overline{\\psi}$, any configuration with a value of $\\psi$ in the negative specific heat branch may be a chance configuration of a cluster that has a negative specific heat.  Since the $6N$-dimensional phase space configuration of subclusters in a cell is directly related to the cell's potential and kinetic energies and thus its virial ratio, we can use the theory to study the internal structure and shapes within a cell. To do this, we relate the energy calculated from the observed instantaneous phase space configuration to the ensemble average energies in quasi-equilibrium. This highlights the fact that clusters of galaxies are not in strict equilibrium and hence their potential and kinetic energies fluctuate about the ensemble average energy. The spatial configurations of clusters change as a result of these fluctuations.  Because the average thermodynamic quantities of an ensemble of cells in quasi-equilibrium change very slowly compared to the dynamical timescale of a single cell, the quasi-equilibrium energies are approximately time averages. We use the spectrum of fluctuations to determine the probability that a cell may have a particular instantaneous virial ratio. This is a necessary but not sufficient condition for a cell to have a given configuration. To compare different configurations, we therefore compare the probability that a cell has a virial ratio which could give rise to a given configuration.  Based on the range of fluctuations in virial ratio observed in $N$-body simulations by \\citet{1972ApJ...172...17A}, we conclude that cells that have a negative specific heat may have almost any spatial configuration. These configurations are likely to be chance configurations of fluctuating cells that are nearly virialized. Cells with positive specific heat have a range of energies corresponding to fewer spatial configurations. In our comparisons between line and ring configurations, the probability that a cell will have a line configuration rather than a ring configuration varies depending on the clustering parameter $b$ and the number $N$ of subclusters in the cell. Therefore we present the following procedure for comparing the shapes of two clusters:  The first step is to define a region of space that covers the cluster. The size and shape of this region of space define the cell size. Within the cell, identify the subclusters whose masses are within an order of magnitude of the largest subcluster. The positions of these subclusters are directly related to the scaled potential energy $W_*$ of the cluster. From the observed $W_*$ or the peculiar velocity information, estimate the virial ratio $\\psi$.  The next step is to use the cell size to calculate the mean and variance of the counts-in-cells distribution of similarly-sized subclusters for a larger sample to obtain the clustering parameter $b$~(e.g. \\citealt{2005ApJ...626..795S,2011ApJ...729..123Y}). Then using the observed parameters $\\psi$ and $b$, calculate the probability of finding a cell with a similar virial ratio using equation \\refeq{eq-PpsiObs}. This probability should then be compared to a reference cell that has a different structure to determine the relative probabilities of different configurations.  Although the configurations that we use here are highly idealized, our comparison of a line and a ring configuration is generally consistent with observations of the cosmic web. This simple test of our theory gives a reasonable result. We intend to make more detailed comparisons with observations and more realistic cases which may include dark matter in a forthcoming paper."
    },
    "1107/1107.4567_arXiv.txt": {
        "abstract": "{Models of chemical evolution of galaxies including the dust are nowadays required to decipher the high-z universe. In a series of three papers we have tackled the problem and set a modern chemical evolution model in which the interstellar medium (ISM) is made of gas and dust. In the first paper (Piovan et al., 2011a) we revised the condensation coefficients for the elements that typically are present in the dust. In the second paper (Piovan et al., 2011b) we have implemented the dust into  the Padova model with infall and radial flows  and tested it against the observational data for the Solar Neighbourhood (SoNe). In this paper (the third of the series) we extend it to the whole Disk of the Milky Way (MW).} {The Disk  is used as a laboratory to analyze the spatial and temporal behaviour of (i) several dust grain families with the aid of which we can describe the ISM dust in a simple way, (ii) the abundances in the gas, dust, and total ISM of the elements present in the dust emitted by stars and the dust grains grown by accretion in the ISM, and finally (iii) the depletion of the same elements.} { The temporal evolution of the dust and gas across the Galactic Disk is calculated under the effect of radial flows  and a Bar in the central regions. The gradients of the abundances of C, N, O, Mg, Si, S, Ca and Fe in gas and dust across the Disk are derived  as a function of time.} {The theoretical gradients nicely reproduce those derived from Cepheids, OB stars, Red Giants and HII regions. This provides the backbone for the companion processes of dust formation and evolution across the Galactic Disk. We examine in detail the contributions to dust by AGB stars, SN{\\ae} and grain accretion in the ISM at different galacto-centric distances across the Disk. Furthermore, we examine the variation of the ratio between silicates and carbonaceous grains with time and position in the Disk. Finally, some hints about the depletion of the elements in regions of high and low SFR (inner and outer Disk) are presented.} {The results obtained for the Disk of the MW make it possible  to extend the model we have developed to other astrophysical situations such as galaxies of different morphological types, like ellipticals or star-busters, or even to a different theoretical modelling, like the chemo-dynamical N-Body simulations.} ",
        "introduction": "\\label{Introduction}  New generation, powerful telescopes and space instrumentation have unveiled a high-z universe surprisingly obscured by large amounts of dust, and consequently spurred a new generation of chemical and spectral galaxy models in which the formation, presence  and evolution of dust is taken into account \\citep[See for instance][to mention the most recent ones]{Zhukovska08,Calura08,Valiante09,Pipino11,Gall11a,Gall11b,Mattsson11,Valiante11,Acharyya11, Kemper11,Dwek11}. Given  the interest in understanding the properties of the early universe, most of these models are tuned toward QSOs or star-burst galaxies. They try to get clues about the role of dust  in the early stages and to estimate how much dust can be produced both by stars and other processes in the interstellar medium (ISM) on a given timescale in order to match the amount of dust observed in high-z objects \\citep{Gall11a,Gall11b,Mattsson11,Dwek11}. However, because of the huge distances, the information on  dust  to our disposal for these extremely far objects comes from (i) the extinction curve, that suggests clues about dust composition by the comparison with the one for galaxies in the  local such as the SMC, LMC, and M; (ii) estimates of the dust mass from the fluxes observed in the pass-bands to our disposal. Other pieces of information are typically related to the upper limit to the age of the target, its metallicity and star formation rate (SFR) and, finally, the total masses in form of  stars, molecular gas (if CO observations are available), and  Dark Matter \\citep{Valiante11}. Since we are observing the integrated properties of these galaxies, the history of chemical enrichment for individual elements and different distances from the centre of the galaxy are hardly available. A detailed information about the way in which the various elements behave is however required if we want to cross compare data with theoretical modelling of dust formation/evolution, because it would correlate the processes affecting the dust with the gas composition existing in that part of the galaxy we intend to model \\citep{Zhukovska08,Piovan11b}. This indeed would be the typical case in which  the relative proportion of silicates/carbonaceous grains and iron dust develop with time \\citep{Dwek05,Zhukovska09} in different regions of a galaxy, thus getting hints on the dust mixture. For its proximity and hence richness of data, the best possible laboratory is the SN. Another useful laboratory in which we can trace the evolution of the element abundances by means of the observations of different types of objects is the Disk of the MW. Nevertheless, till now, only an handful of models have considered the presence of dust in the Galactic Disk and its evolution in space and time \\citep{Dwek98,Zhukovska09}.\\\\  \\indent Basing on the observational data we have referred to above, in this study we plan to examine  the formation/evolution of dust and  depletion of the elements along the MW Disk, trying to satisfy the constraints derived from the radial gradients in  matter and abundances  of  various chemical elements as measured in different types of sources. This paper is the last one of a series of three in which we analyzed the formation and evolution of dust in the MW. It is worth briefly summarizing the aims and results  of the first two papers, and also briefly comment on the gradual development of the chemical model. In the first paper \\citep{Piovan11a} we examined and presented a compilation  of condensation efficiencies to be adopted to study the dust content of the gas. In the second paper \\citep{Piovan11b} we apply our condensation coefficients to a detailed chemical model of the Solar Neighborhood and reproduced the local depletions of the chemical elements together with other observational data getting clues on the effect induced by important parameters such as the IMF, the star formation law, and the accretion model adopted to simulate the grain growth in the ISM. In addition to this, the same analysis has been made for the inner and outer regions of the MW Disk, trying to understand how they would behave in a high/low star forming environment.\\\\  \\indent The plan of the present paper is as follows. In Sect. \\ref{chemical} we summarize the main features of the chemical model, more details can be also found in \\citep{Piovan11b}. In particular, we cast the equations governing the model in presence of radial flow of matter and dust. In Sect. \\ref{discrete} we present the equations for the radial flows in presence of dust, gas, and ISM (gas+dust). In Sect. \\ref{boundaryCond} we discuss the boundary conditions and the final equation for the innermost shell (Sect. \\ref{shell1}), the generic shell (Sect. \\ref{boundary}), and the outermost shell (Sect. \\ref{shellN}). Some computational details are discussed in Sect. \\ref{numericalradf}. In Sect. \\ref{RadialMCs} we present a simple algorithm  based on the Artificial Neural Networks (ANNs) to predict the amount of molecular gas (cold regions where we assume that dust accretion takes place) in the different rings of the Disk taking into account the information about SFR and  surface mass density. In Sect. \\ref{RadialAbundances} we compare the results of the chemical models, whose properties for the Solar Vicinity have been already discussed in \\citet{Piovan11b}, extended now the whole Disk. The theoretical predictions are compared with the observational data for the gradients in gas content, chemical abundances, and abundance ratios. Furthermore, we present and discuss the radial and time evolution of many dust-related quantities. In particular, we analyze the radial behaviour of (i) some typical grain families; (ii) the elements involved in the dust formation, and (iii) the depletion in the innermost and outermost regions. Finally, in Sect. \\ref{Discus_Concl} we discuss the results and draw some general conclusions. ",
        "conclusions": "This is the last paper of a series of three \\citep{Piovan11a,Piovan11b} devoted to the study of star-dust and dust accretion in the ISM. We have analyzed in some detail the chemical gradients (gas and dust) across the Disk of the MW. The MW as a whole and its Disk and SoNe have been considered as ideal laboratories in which theories of chemical evolution and dust formation are compared with observational data. A suitable description of the Disk including radial flows of matter and the presence of a inner Bar has been adopted to simultaneously explain the bump on the inner gas density and the gradients in chemical abundances. This model provides the frame to the formation and evolution of the dust, one of the components of the ISM (the gas elemental species being the others).  The main conclusions of this study can be summarized as follows:  \\begin{itemize} \\item For plausible combinations of the parameters, the same that in \\citet{Piovan11b} allowed to nicely reproduce the properties of the SoNe, both for gas and dust, we find a good general agreement with the mass gradients observed in the MW Disk, both for the total gas mass and the single element abundances. These latter  have been compared with the abundances observed in different types of objects, like Cepheids, open clusters, HII regions, O and B stars and Red Giants. Some striking differences between models and observations for C and N can be easily explained by means of the stellar evolution theories.  \\item The evolutionary scheme for dust formation and evolution across the Disk, emerging from the simulations, is as follows: SN{\\ae} dominate the dust budget in the early stages in the whole disk. As time goes on, AGB stars enter the game: they contribute to the C-based dust budget if the metallicity keeps low,  switch to O-rich dust as soon as the metallicity is near  solar or super-solar. For this reason AGB stars behave in a different way in the inner regions where the metallicity quickly increases and in the outer regions where the metallicity keeps low for  long time. The ISM accretion is the main contributor to the dust budget across the whole disk. It starts very early on  in the innermost regions due to the faster enrichment and the highest number  densities of the elements involved in dust and much later in the outer regions where metallicity and density are lower and where AGB stars of low-Z can be very important for many Gyr. It is clear from this trend that a crucial role is played by the SN{\\ae}: the faster they enrich the ISM in metals, the faster the ISM starts to play a decisive role. Furthermore, the higher  the condensation efficiencies, the higher is the amount of dust that can be injected during the time interval before the ISM dominates. The uncertainties on the efficiency of dust condensation in SN{\\ae} \\citep{Piovan11a} could clearly change the picture during the early evolutionary phases \\citep{Zhukovska08,Gall11a,Gall11b}.  \\item The radial distribution of the dust mass follows the behaviour of the gas showing a bump in the inner zone and a negative slope at the present time. It decreases moving outwards. The fractional radial mass of silicates has also  a \\textit{negative} slope, whereas the carbonaceous grains have a \\textit{positive} slope. They also  play a more and more important role as we move toward the outskirts. This is ultimately due to the combined effect of several causes : (i) the yield from SN{\\ae} that enrich the ISM with elements involved in silicates and dust where star formation is more active; (ii) the long time interval during which low-Z AGB stars inject carbon in the outskirts of the Disk, and (iii) the delay in the onset of the ISM accretion in the outer regions.  \\item The depletion of the elements is more significant in the inner regions, with the highest star formation and densities, while in the outskirts we have less depletion due to lower efficiency of dust accretion and star-dust injection.  \\item As expected, the depletion significantly alters  the radial gradients in the gas  abundances, in particular for the elements involved in silicates, for which (at least in our simplified model) many elements enter the process and the key element driving the process can vary with time, making the whole thing more irregular. \\end{itemize}   We can conclude by saying that the picture drawn by  the chemical model for the Galactic Disk in presence of dust  is  consistent with the way in which SN{\\ae}, AGB stars of different Z and ISM contribute to the total dust budget. In particular, since the radial properties of the elements along the disk are satisfactorily reproduced, we are confident that the dust enrichment scenario described is based upon a consistent general enrichment in metals of the ISM.  \\textit{Acknowledgements}. L. Piovan acknowledges A. Weiss and the Max Planck Institut Fur AstroPhysik (Garching - Germany) for the very warm and friendly hospitality and for providing unlimited computational support during the visits as EARA fellow when a significant part of this study has been carried out. The authors are also deeply grateful to  S. Zhukovska and H. P. Gail for many explanations and clarifications about their model of dust accretion, T. Nozawa and H. Umeda for many fruitful discussions about SNa dust yields. This work has been financed by the University of Padua with the dedicated fellowship \"Numerical Simulations of galaxies (dynamical, chemical and spectrophotometric models), strategies of parallelization in dynamical lagrangian approach, communication cell-to-cell into hierarchical tree codes, algorithms and optimization techniques\" as part of the AACSE Strategic Research Project.   \\begin{scriptsize}"
    },
    "1107/1107.0564_arXiv.txt": {
        "abstract": "{Supernova remnants (SNRs) were often discovered in radio surveys of the Galactic plane. Because of the surface-brightness limit of previous surveys, more faint or confused SNRs await discovery. The Sino-German $\\lambda$6\\ cm Galactic plane survey is a sensitive survey with the potential to detect new low surface-brightness SNRs.} {We want to identify new SNRs from the $\\lambda$6\\ cm survey map of the Galactic plane.} {We searched for new shell-like objects in the $\\lambda$6\\ cm survey maps, and studied their radio emission, polarization, and spectra using the $\\lambda$6\\ cm maps together with the $\\lambda$11\\ cm and $\\lambda$21\\ cm Effelsberg observations. Extended polarized objects with non-thermal spectra were identified as SNRs.} {We have discovered two new, large, faint SNRs, G178.2$-$4.2 and G25.1$-$2.3, both of which show shell structure. G178.2$-$4.2 has a size of $72\\arcmin \\times 62\\arcmin$ with strongly polarized emission being detected along its northern shell.  The spectrum of G178.2$-$4.2 is non-thermal, with an integrated spectral index of $\\alpha = -0.48\\pm0.13$. Its surface brightness is $\\Sigma_{\\rm 1~GHz} = 7.2\\times10^{-23}{\\rm Wm^{-2} Hz^{-1} sr^{-1}}$, which makes G178.2$-$4.2 the second faintest known Galactic SNR. G25.1$-$2.3 is revealed by its strong southern shell which has a size of $80\\arcmin \\times 30\\arcmin$. It has a non-thermal radio spectrum with a spectral index of $\\alpha = -0.49\\pm0.13$.} {Two new large shell-type SNRs have been detected at $\\lambda$6\\ cm in an area of 2200~deg$^{2}$ along the the Galactic plane. This demonstrates that more large and faint SNRs exist, but are very difficult to detect.} ",
        "introduction": "Supernova remnants (SNRs) are the post-explosion relics of massive stars that have reached the end of their evolutionary life times.  It has been predicted that the total number of Galactic SNRs is between 1\\,000 and 10\\,000 \\citep{Berkhuijsen84, Li91, Tammann94}.  However, up to now just 274 SNRs have been identified \\citep{Green09}. Among these, G192.8$-$1.1 has recently been disapproved as being a SNR \\citep{Gao11x}. This low SNR detection rate results from the insufficient sensitivity and resolution of available observations. Diffuse radio emission and discrete source complexes are not uniformly distributed along the Galactic disk, whose own emission confuses that from faint SNRs and hinders their identification.  A large number of SNRs have been identified from radio survey maps \\citep[e.g.][]{Reich88c, Brogan06}. Shell-type SNRs can be recognized by their morphology, non-thermal spectra, and ordered polarization. Most known Galactic SNRs have a bright shell, with intrinsic magnetic fields running along the shell. The radio spectrum of a SNR is often described by single power law, $S_{\\nu} \\sim \\nu^{\\alpha}$. Here $S_{\\nu}$ represents the integrated flux density of a SNR at the observing frequency $\\nu$. In general, SNRs have a spectral index of $\\alpha \\sim -0.5$, which is expected for SNRs in the adiabatic expansion phase with a compression factor of four \\citep[e.g.][]{Reich02}. Young SNRs can have steeper spectra ($\\alpha \\sim \\rm -0.6\\ to -0.8$) and show radial magnetic fields \\citep[e.g.][]{Reich02}.  To discover new SNRs, observations of high sensitivity are needed, which resolve confusing structures. We have conducted the Sino-German $\\lambda$6\\ cm polarization survey of the Galactic plane \\citep{Sun07, Gao10, Sun11, Xiao11} in the region of Galactic longitude $10\\degr \\leq \\ell \\leq 230\\degr$ and latitude $|b| \\leq 5\\degr$, with an angular resolution of 9$\\farcm$5.  The sensitivity of the Urumqi survey, if extrapolated from 4.8~GHz to 1~GHz with a typical spectral index of $\\alpha = -0.5$, is on average $\\Sigma_{\\rm 1~GHz} = 4.9\\times10^{-23} {\\rm W m^{-2} Hz^{-1} sr^{-1}}$, lower than the surface brightness of the faintest Galactic SNR known to date, which is $\\Sigma_{\\rm 1~GHz} = 5.8\\times10^{-23} {\\rm W m^{-2} Hz^{-1} sr^{-1}}$ for G156.2+5.7 \\citep{Reich92, Xu07}. Therefore, it should be possible to detect new faint SNRs in the $\\lambda$6\\ cm survey. Here, we report the discovery of two new SNRs, G178.2$-$4.2 and G25.1$-$2.3, and study their radio properties.  \\begin{figure*} % \\centering \\includegraphics[angle=-90, width=0.32\\textwidth]{g178_6i_1min_color.ps} \\includegraphics[angle=-90, width=0.32\\textwidth]{G178_cl_I_plot_gauss_1.0_color.ps} \\includegraphics[angle=-90, width=0.32\\textwidth]{g178_6pi_1min_color.ps}\\\\[2mm] \\includegraphics[angle=-90, width=0.32\\textwidth]{g178_11i_1min_color.ps} \\includegraphics[angle=-90, width=0.32\\textwidth]{G178_cl_I_plot_gauss_1.0_fchop_color.ps} \\includegraphics[angle=-90, width=0.32\\textwidth]{g178_11pi_1min_color.ps}\\\\[2mm] \\includegraphics[angle=-90, width=0.32\\textwidth]{g178_21i_1min_color.ps} \\includegraphics[angle=-90, width=0.32\\textwidth]{G178_21_cl_I_plot_gauss_1.0_color.ps} \\includegraphics[angle=-90, width=0.32\\textwidth]{g178_21pi_1min_color.ps}\\\\ \\caption{ Radio images of G178.2$-$4.2 at $\\lambda$6\\ cm from the Sino-German $\\lambda$6\\ cm polarization survey in the {\\it top panels}, at $\\lambda$11\\ cm which we newly observed in the {\\it middle panels}, and at $\\lambda$21\\ cm from the Effelsberg Medium Latitude Survey (EMLS) \\citep{Reich04} in the {\\it bottom panels}. The angular resolutions of the $\\lambda$6\\ cm, $\\lambda$11\\ cm and $\\lambda$21\\ cm maps are 9.5$\\arcmin$, 4.4$\\arcmin$ and 9.4$\\arcmin$, respectively. The {\\it left panels} show the total-intensity ($I$) maps in color and in contours, with observed $B$-vectors overlaid (i.e. the observed $E$-vectors plus $90\\degr$) for polarization intensities $PI > 2.4~{\\rm mK}\\ T_{B}$ at $\\lambda$6\\ cm and 32.0~${\\rm mK}\\ T_{B}$ at $\\lambda$11\\ cm. Vectors are not shown at $\\lambda$21\\ cm. The vector length is proportional to $PI$. The $I$ contours in the $\\lambda$6\\ cm maps are $2^{\\frac{n-1}{2}}\\times1.6~(4\\sigma)~{\\rm mK}\\ T_{B} $, (n = 1, 2, 3 ...), in the $\\lambda$11\\ cm maps are $2^{\\frac{n-1}{2}}\\times12.8~ (4\\sigma)~{\\rm mK}\\ T_{B}$ (n = 1, 2, 3 ...), and in the $\\lambda$21\\ cm maps are $2^{\\frac{n-1}{2}}\\times44.0~(4\\sigma)~{\\rm mK}\\ T_{B} $ (n = 1, 2, 3 ...). The {\\it central panels} also display the total-intensity maps, where strong point-like sources have been subtracted, to show the extended emission from the SNR more clearly. The {\\it right panels} are the polarization intensity images with the contours for the total intensity maps with point-like sources subtracted. The rectangle in the $\\lambda$6\\ cm image in the {\\it top central panel} outlines the area used for the TT-plot spectral analysis displayed in Fig.~\\ref{tt}.} \\label{g178_maps} \\end{figure*}  \\begin{figure*} \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_i_1min_color.ps} \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_cl_1min_color.ps} \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_pi_1min_color.ps}\\\\[2mm] \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_11i_1min_color.ps} \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_11icl_1min_color.ps} \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_11pi_1min_color.ps}\\\\[2mm] \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_21i_1min_color.ps} \\includegraphics[angle=-90, width=0.325\\textwidth]{g25_i21cl_1min_color.ps} \\caption{Same as Fig.~\\ref{g178_maps}, but for G25.1$-$2.3. Here the $\\lambda$11\\ cm and $\\lambda$21\\ cm maps were extracted from the Effelsberg $\\lambda$11\\ cm survey \\citep{Reich9011} and $\\lambda$21\\ cm survey \\citep{Reich9021}, respectively. The contours in the $\\lambda$6\\ cm total-intensity maps are $ 2^{\\frac{n-1}{2}}\\times12.8~(4\\sigma)~{\\rm mK}\\ T_{B} $ (n = 1, 2, 3 ...), in the $\\lambda$11\\ cm total-intensity maps are $2^{\\frac{n-1}{2}}\\times72.0~(4\\sigma)~{\\rm mK}\\ T_{B} $ (n =1, 2, 3 ...), and in the $\\lambda$21\\ cm total-intensity maps are $2^{\\frac{n-1}{2}}\\times200.0~(4\\sigma)~{\\rm mK}\\ T_{B} $ (n = 1, 2, 3 ...). The polarization intensity threshold for the $B$-vectors is 12.0~${\\rm mK}\\ T_{B}$ at $\\lambda$6\\ cm and ${\\rm 75~mK}\\ T_{B}$ at $\\lambda$11\\ cm. No polarization data at $\\lambda$21\\ cm are presently available. The star-symbols in the $\\lambda$6\\ cm maps ({\\it top left panel}) indicate known pulsars in the field. The polygon in the $\\lambda$6\\ cm image in {\\it top central panel} outlines the area for the TT-plot spectral analysis shown in Fig.~\\ref{tt}.} \\label{g25_maps} \\end{figure*} ",
        "conclusions": "We have found two shell-like objects, G178.2$-$4.2 and G25.1$-$2.3, from the radio map of the Sino-German $\\lambda$6\\ cm polarization survey of the Galactic plane. In addition, using the Effelsberg $\\lambda$11\\ cm and $\\lambda$21\\ cm continuum and polarization maps, a shell-type morphology is confirmed for both. Their radio spectra are characteristic of non-thermal emission with spectral indices for the shells of $\\alpha \\sim -0.6$ for G178.2$-$4.2 and $\\alpha \\sim -0.5$ for G25.1$-$2.3. Such values are typical for SNRs.  An ordered magnetic field runs along the northern shell of G178.2$-$4.2. An \\ion{H}{I} cavity, likely at a distance of about 2.9~kpc, is probably associated with the SNR G25.1$-$2.3. This would imply a diameter of up to 67~pc in size."
    },
    "1107/1107.1117.txt": {
        "abstract": "Primordial or Big Bang nucleosynthesis (BBN)  is one of the three strong evidences for the Big-Bang model together with the expansion of the Universe and the Cosmic Microwave Background radiation. In this study, we improve the standard BBN calculations taking into account new nuclear physics analyses  and we enlarge the nuclear network until Sodium. This is, in particular,  important to evaluate the primitive value of CNO mass fraction that could affect Population III stellar evolution. For the first time we list the complete network of more than 400 reactions with references to the origin of the rates, including $\\approx$270 reaction rates calculated using the TALYS code. Together with the cosmological light elements, we calculate the primordial Beryllium, Boron, Carbon, Nitrogen and Oxygen nuclei. We performed a sensitivity study to identify the important reactions  for  CNO, \\neu\\ and Boron nucleosynthesis. We reevaluated those important reaction rates using experimental data and/or theoretical evaluations. The results are compared with precedent calculations: a primordial Beryllium abundance increase by a factor of 4 compared to its previous evaluation, but we note a stability for B/H and for the CNO/H abundance ratio that remains close to its previous value of  $0.7\\times10^{-15}$. On the other hand, the extension of the nuclear network has not changed the \\sep\\ value, so its abundance is still  3--4 times greater than its observed spectroscopic value. ",
        "introduction": "There are presently three observational evidences for the Big-Bang Model: the universal expansion, the Cosmic Microwave Background (CMB) radiation and Primordial or Big-Bang Nucleosynthesis (BBN). The third evidence for a hot Big-Bang comes from the primordial abundances of the ``light elements'': \\hli. They are produced during the first $\\approx$20 min of the Universe when it was dense and hot enough for nuclear reactions to take place.   The number of free parameters entering in Standard BBN has decreased with time. The number of light neutrino families is known from the measurement of the $Z^0$ width by LEP experiments at CERN: $N_\\nu$ = 2.9840$\\pm$0.0082~\\citep{LEP}. The lifetime of the neutron entering in weak reaction rate calculations and many nuclear reaction rates have been measured in nuclear physics laboratories. The last parameter to have been independently determined is the baryonic density of the Universe which is now deduced from the observations of the anisotropies of the CMB radiation coming from the WMAP satellite. The number of baryons per photon  which remains constant during the expansion,  $\\eta$, is directly related to \\ob\\ by \\obh=3.65$\\times10^7\\eta$ . The WMAP7 gives now : \\obh=0.02249$\\pm$0.00056  and $\\eta$ = 6.16 $\\pm$ 0.15 $\\times 10^{-10}$ \\citep{WMAP}. In this context, primordial nucleosynthesis is a parameter free theory and is the earliest probe of the Universe. We note an overall agreement except for the \\sep. In the literature many studies have been devoted to this lithium problem \\citep[and references therein]{Ang05, Coc07, Cyburtetal2008, Cha11} .   Deuterium, a very fragile isotope, is destroyed after BBN. Its most primitive abundance is determined from the observation of clouds at high redshift, on the line of sight of distant quasars. Very few observations of these cosmological clouds  are available \\citep[and references therein]{pettini08} and the adopted primordial D abundance is given by the average value:  D/H~$ = (2.82^{+0.20}_{-0.19})\\times 10^{-5}$.     After BBN, \\qua\\ is still produced by stars. Its primitive abundance is deduced from observations in HII (ionized hydrogen) regions of compact blue galaxies. Galaxies are thought to be formed by the agglomeration of such dwarf galaxies which are hence considered as more primitive. The primordial \\qua\\ abundance $Y_p$ (mass fraction) is given by the extrapolation to zero metallicity but is affected by systematic uncertainties \\citep{aver10, isot10} such as plasma temperature or stellar absorption. These  most recent determinations based on almost the same set of observations lead to:  $Y_p = 0.2561 \\pm 0.0108$.   Contrary to \\qua\\ ,  \\tro\\ is both produced and destroyed in stars so that the evolution of its abundance as a function is subject to large uncertainties and has only been observed in our Galaxy \\citep{Ban02},  \\tro\\ /H~$ = 1.1 \\pm 0.2 \\times 10^{-5}$.  \\noindent Consequently, the baryometric status of \\tro\\ is not firmly established \\citep{vang03}.  Primordial lithium abundance is deduced from observations of low metallicity stars in the halo of our Galaxy where the lithium abundance is almost independent of metallicity, displaying a plateau, the so-called Spite plateau \\citep{Spite82}. This interpretation assumes that lithium has not been depleted at the surface of these stars, so that the presently observed abundance is supposed to be equal to the initial one. The small scatter of values around the \u00d2Spite plateau\u00d3 is an indication that depletion may not have been very effective.  Astronomical observations of these metal poor halo stars \\citep{Ryanetal2000} have led to a relative primordial abundance of:  Li/H~$ = (1.23^{+0.34}_{-0.16}) \\times 10^{-10}$ .  \\noindent A more recent analysis by \\citet{sbordone10} gives:  Li/H~$ =  (1.58 \\pm 0.31) \\times 10^{-10}$.  \\noindent  More generally, \\citet{Spite10} have reviewed the last Li observations and their different astrophysical aspects. See also \\citet{frebel11} for a wide review.    In 2006, high-resolution observations of Li absorption lines in some very old halo stars have also claimed evidence for a large primitive abundance  of the weakly-bound isotope $^6$Li \\citep{Asp06}. The \\six/\\sep\\ ratios of $\\sim 5\\times 10^{-2}$ were found to be about three orders of magnitude larger than the BBN-calculated value of \\six/\\sep$\\sim 10^{-5}$. The key BBN \\six\\ production mechanism is the D($\\alpha$,$\\gamma$)$^6$Li reaction at energies in the range of 50~keV $\\le E_{cm} \\le$ 400~keV~\\citep{Ser04}. This reaction has very recently been re-investigated \\citep{Ham10} confirming the previous result that \\sbbn\\ cannot produce \\six\\ at the required level. Concerning the \\six\\ observational status, more recently, however, \\citet{Cay07} and \\citet{Ste10} have pointed out that line asymmetries similar to those created by a $^6$Li blend could also be produced by convective Doppler shifts in stellar atmospheres. In this context,  these observations have to be confirmed. More detailed analyses are necessary to firmly conclude about the detection of the \\six\\ abundance at this level in these metal poor stars \\citep[see][for a review]{Spite10}.   We consider in detail in this present study the other isotopes potentially produced by the \\sbbn\\ including \\neu, \\dix, \\onz\\ and the CNO isotopes. The production of \\neu, \\dix\\ and \\onz\\ (BeB) and CNO isotopes have been studied in the context of standard and inhomogeneous BBN \\citep{Tho93, Tho94, kajino1, kajino2, kajino3, Ioc07, Ioc09}. The most relevant analysis concerning CNO in BBN comes from  \\citet{Ioc07} who included more than 100 nuclear reactions and predicted a CNO/H abundance ratio  of approximately $6 \\times 10^{-16}$, with an upper limit of  $10^{-10}$. These evaluations had been also performed to provide the initial conditions for the evolution of population III stars.  BeB nucleosynthesis is an important chapter of nuclear astrophysics. Specifically, rare and fragile nuclei, they are not generated in the normal course of stellar nucleosynthesis and are, in fact, destroyed in stellar interiors. This characteristic is reflected by the low abundance observations of these light  species \\citep[and references therein]{prim10, boesg10}. A glance to the abundance curve  suffices to capture the essence of the problem: a gap separates He and C. At the bottom of this precipice rests the trio Li-Be-B. At very low metallicity, the BeB abundance is less than $10^{-12}$ relatively to hydrogen. Indeed, they are characterized by the simplicity of their nuclear structure (6 to 11 nucleons) and their scarcity in the Solar System and in stars. In fact, they are fragile because a selection principle at the nuclear level has operated in nature. Due to the fact that nuclei with mass  5 and 8  are unstable, BBN has almost stopped at A = 7, while nuclear burning in stars bypasses them through the triple--alpha reaction.  After BBN, the formation agents of LiBeB are Cosmic Rays  interacting with interstellar or circumstellar CNO. Other possible origins have been also identified, for example supernova neutrino spallation (for \\sep\\ and \\onz\\ ). In contrast, \\six, \\neu\\ and \\dix\\ are pure cosmic-ray spallative products.  (For a review see \\citet{vang00}.)  Recently, the LiBeB production has been considered in a cosmological context \\citep{rollinde06, rollinde08}. The non-thermal evolution with redshift of \\six, Be, and B in the first structures of the Universe has been studied. In this context  cosmic rays are impinging alpha particles and CNO produced by the first massive stars, the so-called Population III stars. The computation has been  performed in the framework of hierarchical structure formation and reliable \\six\\ and BeB initial abundances coming from BBN are required to optimize the initial conditions.   Even though the direct detection of primordial CNO isotopes seems highly unlikely with the present observational techniques at high redshift, it is also important to better estimate their \\sbbn\\ production. Hydrogen burning in the first generation of stars (Pop III stars) proceeds through the slow pp chains until enough carbon is produced (through the triple-alpha reaction) to activate the CNO cycle. The minimum value of the initial CNO mass fraction that would affect Pop III stellar evolution is estimated to be 10$^{-10}$ \\citep{Cas93}  or even as low as  10$^{-12}$ in mass fraction for the less massive ones \\citep{Eks08}. This is only two orders of magnitude above the  \\sbbn\\ CNO yield using the current nuclear reaction rate evaluations of \\citet{Ioc07}.  In addition, it has been shown that Pop III stars evolution is sensitive to the triple-alpha $^{12}$C producing reaction and can be used to constrain the possible variation of the fundamental constants \\citep{Eks10}. This reaction is sensitive to the position of the Hoyle state, which in turn is sensitive to the values of the fundamental constants. The amount of produced CNO ($^{12}$C) could affect the HR diagram (CNO versus pp H-burning) and the final production of $^{12}$C and $^{16}$O in Pop III stars. Hence, it is important to quantify the amount of primordial CNO present at their birth. In the same context of the variations of the fundamental constants, $^8$Be (which decays to two alpha particles within $\\sim10^{-16}$~s) could become stable if these constants were only slightly different. At BBN time, this would possibly allow to bridge the \"A=8 gap\" and produce excess CNO. To determine how significant would be this excess, one needs to know the standard  BBN production of the CNO elements.  Another motivation for this study is the above mentioned dichotomy concerning the Li abundance. At WMAP baryonic density, \\sep\\ is produced as $^7$Be that later decays. Nuclear ways to destroy this $^7$Be have been explored. An increased $^7$Be(d,p)$2\\alpha$ cross section has been proposed by \\citet{Coc04} but was not confirmed by experiment described in  \\citet{Ang05} unless a new resonance is present \\citep{Cyb09} with very peculiar properties. Other $^7$Be destruction channels have recently been proposed by \\citet{Cha11} awaiting experimental investigation. Another scenario would be to take advantage of an increased late time neutron abundance. This is exactly what happens (in the context of varying constants)  when the $^1$H(n,$\\gamma)^2$H rate is decreased. The neutron late time abundance is increased (with no effect on \\qua) so that more $^7$Be is destroyed by  $^7$Be(n,p)$^7$Li(p,$\\alpha)\\alpha$.  \\citep[see in] [Fig. 1]{Coc07}.  The main difficulty in BBN calculations up to CNO is the extensive network needed, including n-, p-, $\\alpha$-, but also d-, t- and  \\tro-induced reactions. Most of the corresponding cross sections cannot be extracted from experimental data only. This is especially true for radioactive tritium-induced reactions, or for those involving radioactive targets like e.g. $^{10}$Be. For some reactions, experimental data, including spectroscopic data of the compound nuclei, are just inexistent. Hence, for many reactions, one has to rely on theory to estimate the reaction rates. Previous studies \\citep{Tho93,Tho94,Ioc07} have performed unpublished analyses of experimental data but have also apparently extensively used the prescription of \\citet{Fow64} and \\citet{Wag67} to estimate many rates. These prescriptions often assume a constant astrophysical $S$-factor. This obviously cannot be a good approximation for most of the reactions considered in the BBN context. In this study we use at first more reliable rate estimates provided by the TALYS reaction code \\citep{TALYS}, next we perform a sensitivity study and, finally improve the rate estimates of the most important reactions by dedicated evaluations.   A detailed analysis of all reaction rates and associated uncertainties would be desirable but is unpractical for a network of $\\approx$400 reactions. So, in Section 2 we present the extended network and the standard thermonuclear reaction rates used. In Section 3, we study the sensitivity of the calculated primordial abundances up to CNO to variation of the rates by a factor of up to 1000 and re-evaluate selected reaction rates. In Section 4, we present the BBN calculation results  and we conclude in section 5. Note that in the annex we give the full list of reactions with the references at the origin of the reaction rates. ",
        "conclusions": "We have used an extensive network of more than 400 nuclear reactions whose thermonuclear reaction rates are adopted from evaluations based on experimental data or were calculated using the TALYS code. It enabled us to calculate the \\sbbn\\ production of \\six, \\neu, \\dix, \\onz\\ and CNO isotopes  more reliably. We performed a sensitivity study by varying uncertain reaction rates by factors of up to 1000 and down to 0.001. In that way, a few reactions that could affect the A$>$7 isotope yields were identified and their rate (re-)evaluated using available experimental data and/or theoretical or phenomenological input. On the basis of these new evaluations the  \\neu, \\dix, \\onz\\ and CNO isotope production was found to be close to the initial results and in global agreement with previous calculations \\citep{Ioc07}.   For most of the few reactions which were identified to have an impact on \\sbbn, we were able to collect sufficient experimental data to derive new reaction rates with associated uncertainties much reduced with respect to our initial three orders of magnitude variation. (An exception is $^7$Be(t,p)$^9$Be but affecting the  \\neu\\ production only.) In some cases these new rates differ from the previous ones by large factors but changes compensating each other (e.g. $^{11}$B(d,n)$^{12}$C and $^{11}$B(d,p)$^{12}$B) allow us to confirm the CNO \\sbbn\\ production of about $0.7\\times10^{-14}$ in mass fraction (i.e CNO/H $\\approx0.7\\times10^{-15}$ ). Based on the rate uncertainties (a factor of $\\lap$10 at BBN temperatures) obtained in Section~\\ref{s:improved} and combining \\citep[see][]{ILCCF10a} the corresponding uncertainty factors from Table~\\ref{t:sensib} we can estimate the uncertainty on CNO production to be of a factor of  $\\lap$4. This is too small to have presently an impact on Population III stellar evolution. It is nevertheless a reference value for comparison with non-\\sbbn\\ CNO production e.g. in the context of varying constants.   Finally, our extension of the network does not help alleviate the discrepancy between the calculated and observed  \\sep\\  abundances. No new late-time neutron source was found that could destroy \\be\\ before it decays to \\sep.  Nevertheless we pointed out the unexpected high sensitivity of the CNO abundance with respect to the $^7$Li(d,n)2$^4$He reaction rate. A similar situation was found between the \\sep\\ abundance  and the $^1$H(n,$\\gamma)^2$H reaction rate. This emphasizes again the complex nature of nucleosynthesis and a posteriori justifies such a sensitivity study: the impact of a given reaction being not always predictable, even in the simple case (i.e. homogeneity, no mixing, nor convection,...) of BBN. As a consequence, even though very unlikely, the search for a nuclear solution to the lithium problem remains justified."
    },
    "1107/1107.5207_arXiv.txt": {
        "abstract": "{\\em Kepler Mission} results are rapidly contributing to fundamentally new discoveries in both the exoplanet and asteroseismology fields.  The data returned from {\\em Kepler} are unique in terms of the number of stars observed, precision of photometry for time series observations, and the temporal extent of high duty cycle observations.  As the first mission to provide extensive time series measurements on thousands of stars over months to years at a level hitherto possible only for the Sun, the results from {\\em Kepler} will vastly increase our knowledge of stellar variability for quiet solar-type stars. Here we report on the stellar noise inferred on the timescale of a few hours of most interest for detection of exoplanets via transits. By design the data from moderately bright {\\em Kepler} stars are expected to have roughly comparable levels of noise intrinsic to the stars and arising from a combination of fundamental limitations such as Poisson statistics and any instrument noise.  The noise levels attained by {\\em Kepler} on-orbit exceed by some 50\\% the target levels for solar-type, quiet stars. We provide a decomposition of observed noise for an ensemble of 12th magnitude stars arising from fundamental terms (Poisson and readout noise), added noise due to the instrument and that intrinsic to the stars.  The largest factor in the modestly higher than anticipated noise follows from intrinsic stellar noise. We show that using stellar parameters from galactic stellar synthesis models, and projections to stellar rotation, activity and hence noise levels reproduces the primary intrinsic stellar noise features. ",
        "introduction": "As of early 2011 the {\\em Kepler Mission} is roughly half way through its baseline 3.5 year mission, with about one-third of the baseline data having now received some analysis. The pace of discoveries and results from the {\\em Kepler Mission} are accelerating and already quite extensive, even though the prime goal of detecting true Earth-analog exoplanets remains in the future, simply from needing to have more extensive time coverage in order to detect transits spaced by a full year. In the area of exoplanets, over 1200 candidates have been detected and reported in \\citet{bor11}, with validation of exoplanets first following from transit timing variations discussed by \\citet{hol10}, the smallest rocky planet to date in \\citet{bat11}, and a six-planet system unveiled in \\citet{lis11}. Discoveries from {\\em Kepler's} second area of emphasis are equally impressive ranging from detection of oscillations on solar-type stars for hundreds of cases in \\citet{cha11a} (only tens had been known before the {\\em Kepler Mission}) to detailed inferences on individual stars (\\citealp{met10,kur11}), and a recent breakthrough involving the detection of both the expected p-modes as well as g-modes allowing sensitive new studies of red giant stars (\\citealp{bed11,bec11}). In stellar astrophysics {\\em Kepler} has revealed a remarkable eccentric A-star binary viewed nearly face on with reflected light features superimposed on tidally-forced g-mode oscillations \\citep{wel11}, and two instances of mutually eclipsing hierarchical triple systems (\\citealp{car11,der11}).  Returning to a discussion of the prime goal of detecting true Earth-analogs, the observational quest is truly daunting.  An Earth-analog presents a transit signature with a depth of 85 parts per million (ppm), lasting a statistical average of 10 hours, and occurring once per year.  The signal would only be evident after seeing three successive transits equally spaced in time. Since the geometric probability of having transits in Earth-analog systems is only 0.5\\%, it is necessary to monitor a very large number of stars nearly continuously over multiple years at very high precision to have any chance of success.  The {\\em Kepler Mission} was scoped \\citep{koc10} to be able to determine the frequency of Earth-size planets at sufficiently long orbital periods to reach well into the habitable zone \\citep{kas93}.  This is challenging on a number of levels, including perhaps the simplest one of the noise levels needed to support detection.  Notable in the \\citet{bor11} paper is a significant fall-off in planet frequency for the smallest planet candidates, especially at the long orbital periods associated with habitability (possible presence of liquid water on the surface based on equilibrium temperatures).  Some of this fall-off is expected to date, since insufficient time has passed to be able to detect true Earth analogs. However, the relative sparsity of Earth-sized planets at shorter orbital periods is open to interpretation and will surely be the focus of much discussion in the exoplanet community.  Do the statistical trends primarily follow from the physics of planet formation (\\citealp{gou11,ida04}), or is there a strong contribution from observational incompleteness given the larger than anticipated noise levels (\\citealp{jen10b,chr10})?  This question is beyond the scope of this paper to resolve.  We will endeavor to more fully explore the noise levels obtained in the {\\em Kepler Mission} given that these are not as low as expected.  The approach in the remainder of the paper is as follows. A self-contained introduction to the mission and data characteristics will be given in \\S 2. Based on consideration of the {\\em Kepler} data set itself, and knowledge of instrument characteristics such as detector readout noise, a subset of $\\sim$12th magnitude stars will be used to establish a decomposition of noise into temporal, instrument, and intrinsic stellar terms in \\S 3. \\S 4 will provide simulations of intrinsic stellar noise as a means of assessing the reasonableness of the noise separation, and for motivating what may be learned about stellar astrophysics from the {\\em Kepler} results. A fainter sample of stars will be reported on in \\S5. Prospects for confirmation of the stellar noise differences found will be touched upon in \\S 6 with a summary in \\S 7. ",
        "conclusions": "We have shown that the noise levels resulting for {\\em Kepler} observations can be decomposed into a few terms:  basics such as Poisson statistics and readout noise, an instrument term that depends on channel number over the 84 amplifiers, a temporal term that depends on specific observing conditions encountered during individual Quarters of observation, and intrinsic stellar noise. The dominant term for roughly solar-type 12th magnitude stars in the overall noise budget is found to follow from the stars themselves.  Excess instrument noise does exist, but is more-or-less in line with expectations, and contributes little to the overall noise within which {\\em Kepler} planet searches must be conducted. By contrast the intrinsic stellar noise, although still very modest at less than 20 ppm for the highest concentration of stars, is a factor of two larger than had been budgeted for. This results in CDPP estimates for {\\em Kp} $\\sim$ 12 that are 50\\% larger than anticipated.  We have shown via simulations of expected fundamental stellar parameters over the {\\em Kepler} field of view, followed by projections of stellar rotation and resulting activity levels, that stellar variability consistent with that observed can be reproduced to first order. These simulations at {\\em Kp} $\\sim$12 produce a broad distribution of stars with intrinsic noise levels over $\\sim$10 -- 20 ppm that is consistent with that derived directly from the {\\em Kepler} data.  CoRoT provided fundamental advances in time-series photometry securely establishing that most, if not all red giants are variable and reaching impressive new precision levels for 12th magnitude dwarfs \\citep{aig09} sufficient to allow searches for small planet transits. {\\em Kepler} has taken this a large step further and is the first mission capable of quantifying the variability of large numbers of stars to the small levels by which the Sun is known to vary. {\\em Kepler} will continue to provide exciting new insights into the astrophysics of quiet stars, and their galactic distributions. While we are not surprised to have learned new things from this new observational capability, the fact that the stars are more variable than expected has a significant influence on the ability to readily detect Earth-analog planet transits where the expected signal per transit is only a few times the inferred noise level on comparable time scales. Observing for a longer time baseline can compensate for the loss of transit detection sensitivity from the higher than anticipated stellar noise."
    },
    "1107/1107.3567_arXiv.txt": {
        "abstract": "We describe a novel multi-layered metal mesh achromatic half wave plate for use in astronomical polarimetric instruments. The half wave plate is designed to operate across the frequency range from 125--250\\,GHz. The wave plate is manufactured from 12-layers of thin film metallic inductive and capacitive grids patterned onto polypropylene sheets, which are then bonded together using a hot pressing technique. Transmission line modelling and 3-D electromagnetic simulations are used to optimize the parameters of the metal-mesh patterns and to evaluate their optical properties. A prototype half wave plate has been fabricated and its performance characterized in a polarizing Fourier transform spectrometer. The device performance is consistent with the modelling although the measured differential phase shift for two orthogonal polarizations is lower than expected. This difference is likely to result from imperfect patterning of individual layers and misalignment of the grids during manufacture. ",
        "introduction": "A wave plate is an optical component which imparts a differential phase shift to light passing through it in different linear polarization directions. This changes the polarization state of a transmitted wave according to the magnitude of this angular phase difference. A half wave plate (HWP) imparts a 180$^{\\circ}$ differential phase shift between light with orthogonal electric field components and can be used to rotate the polarization vector by an angle, $2\\theta$ by rotating the HWP by an angle $\\theta$ with respect to the incident radiation. HWPs are used as polarization modulators in astronomical instruments at millimeter wavelengths ($\\lambda$ = 1--3\\,mm), in particular for measurements of the residual polarization in Cosmic Microwave Background (CMB) anisotropies \\cite{Savini2006,Pisano2006}. These measurements require large spectral bandwidths, cryogenic detectors, and polarimeters with low instrumental polarization and low instrumental emission to ensure photon noise limited detection of the weak polarized CMB anisotropies. This has motivated the development of large area achromatic HWPs \\cite{Savini2006,Pisano2006}.  Conventional HWP technologies use birefringent crystalline materials such as Quartz or Sapphire to provide the differential phase delay of orthogonal polarization components. While single crystal plates demonstrate good performance in a narrow frequency band, a broad-band design was reported initially for use at optical wavelengths \\cite{Pancharatnam1955,Rosenberg1981}. This design incorporates 3 or 5 crystal plates of the same thickness, but differentially rotated with respect to the incoming polarized beam to provide a wider range of wavelengths over which the retardation between orthogonal beams is constant. The first achromatic multiple plate HWPs fabricated for a millimeter wave CMB experiment \\cite{Ade2008} used a sapphire 5-plate design to cover the frequency range from 85--185\\,GHz \\cite{Savini2006,Pisano2006}. This design has been shown to perform well but is limited in aperture size (290\\,mm for sapphire \\cite{Matsumura2009}) by the availability of materials. Furthermore, the required HWP thickness ($\\approx$17\\,mm) leads to a small absorption loss (3\\,\\%) which becomes a significant source of radiative power and therefore photon noise compared to the 2.73\\,K CMB, if the HWP is used at 300\\,K. Lowering the temperature to $\\approx$100\\,K eliminates this emission but significantly complicates the instrument design and cost.  To overcome both the limited diameter and the absorption losses in crystalline wave plates, an alternative method of making HWPs based on metal-mesh technology has been proposed \\cite{Shatrow1995,Pisano2008}. This design uses capacitive and inductive metal-mesh geometries which each generate a frequency dependent phase shift with opposite sign. A first device fabricated from 6 capacitive grids and 6 inductive separated by air/vacuum gaps using accurately etched annular spacers demonstrated a bandwidth of 40\\,\\% (60\\,GHz) with a center frequency of 150\\,GHz \\cite{Pisano2008}. Here we describe a procedure for optimizing the design of this type of HWP and present measurements of a device fabricated with a new process using dielectric spacers between the metallized sheets fused together with a hot pressing technique to make a solid self-supporting disc \\cite{Zhang2009}. This leads to a robust and easily formed component with small absorption losses and a diameter limited by the available size of the patterned meshes which can be manufactured with diameters $\\geq$300\\,mm \\cite{Ade2006}.  In Section 2 of this paper we describe the theory of operation of the HWP and the optimization procedure. We also present the design parameters and expected performance of a HWP optimized for the 125--250\\,GHz region. In section 3 we describe the device manufacture and in section 4 we present optical performance measurements from its characterization as a HWP using a polarizing FTS. ",
        "conclusions": "It is important to put the performance of this HWP into context with the current crystalline HWPs. The two represent very different technologies and both have manufacture and performance issues. Here we attempt to summarize these differences to aid in their selection for future instruments.  Although the prototype has been designed and built at mm-wavelengths, its broadband nature and impedance characteristics are conserved when scaling the geometries to operate at higher frequencies (at sub-mm wavelengths and frequencies higher than 1\\,THz), where crystalline absorption is indeed a problem. A return loss of about ten percent is experienced for this prototype metal mesh HWP without applying additional anti-reflection coating; its average transmission is an improvement with respect to most crystalline birefringent waveplates which have a large impedance mismatch due to the high index of refraction of the crystal material. Single layer anti-reflection coatings, while improving the insertion loss of the crystal alternatives, cannot avoid the inherent crystalline absorption, which gets progressively worse at higher frequencies \\cite{Lau2006}. While reflections can present a problem due to ghost images appearing in a polarimeter system which has more than one reflecting element, absorption is a more serious problem for low background sub-mm astronomical instruments, since the detectors employed in these are background photon noise limited. Thermal grey body emission from a room temperature HWP, which for low loss is proportional to its absorption, can present a significant additional noise component. Most experiments are thus forced to position the HWP modulator inside the cryogenic housing at temperatures below 100K, since the absorption decreases and the grey body emission is also reduced.  While cryogenic deployment of the hot-pressed sample in question, which is unnecessary, would not constitute a problem, in the case of AR coated crystals large temperature gradients require great attention and planning in both the construction process and in the mount design.  Although the recovered phase is not an improvement on existing achromatic crystalline devices \\cite{Savini2006, Savini2009}, the modulation efficiency of the HWP is always above 85\\,\\% in a 90\\,\\% spectral bandwidth with the best performance above 96\\,\\% modulation efficiency in two 10\\,\\% narrow bands (centered at 125 and 295\\,GHz) at the extremes of its spectral range.  In conclusion we note that overall there are numerous instrumental advantages of this device that make it a competitive alternative to crystalline HWP, especially at higher frequencies and where wide spectral bandwidth is required."
    },
    "1107/1107.1879_arXiv.txt": {
        "abstract": "The broad-band spectral distributions of non-thermal sources, such as those of several known blazars, are well described by a log-parabolic fit. The second degree term in these fits measures the curvature in the spectrum. In this paper we investigate whether the curvature parameter observed in the spectra of the synchrotron emission can be used as a fingerprint of  stochastic acceleration.  As a first  approach we use the multiplicative Central Limit theorem to show how fluctuations in the energy gain result in the broadening of the spectral shape, introducing a curvature into the energy distribution.  Then, by means of a Monte-Carlo description, we investigate how the curvature produced in the electron  distribution is linked to the diffusion in  momentum space. To get a more generic  description of the problem we turn to the diffusion equation in  momentum space. We first study some ``standard'' scenarios, in order to understand the conditions that make the curvature in the spectra significant, and the relevance of cooling during the acceleration process. We try to quantify the correlation between the curvature and the diffusive process in the pre-equilibrium stage, and investigate how the transition between the Klein-Nishina  and the Thompson regime, in Inverse Compton cooling, determine the curvature in the distribution at  equilibrium. We apply these results to some observed trends, such as the anticorrelation between the peak energy and the curvature term observed in the spectra of Mrk 421, and a sample of BL Lac objects whose synchrotron emission peaks at X-ray energies ",
        "introduction": "\\label{sec:Intro}  A defining feature of the non-thermal emission from different types of galactic and extragalactic sources is that their spectra are described by a power-law (PL) over a broad photon energy range.  In several sources, however, their spectra show  significant curvature that is typically milder than that expected from an exponential cut-off.  In previous papers \\cite{Massaro2004,Massaro2006}  discussed the curvature observed in the broad band X-ray spectra of the two well known HBL (High-energy peaked BL Lac) objects Mkn~421 and Mkn~501. The basic idea was that this curvature was not simply the result of radiative cooling of high energy electrons, responsible of the synchrotron and inverse Compton emission, but that it was essentially related to the acceleration mechanism. \\cite{Massaro2004}  showed that  curved spectral distributions, in particular log-parabolic (i.e. log-normal) ones, develop when the acceleration probability is a decreasing function of the electron energy. In subsequent works, through the analysis of a large collection of X-ray observations of Mkn 421, \\cite{Tramacere2007Mrk421,TramacerePhD2007,Tramacere2009} pointed out that the observed anticorrelation between the peak energy and the curvature measured in the synchrotron Spectral Energy Distribution (SED), could be used as a clear signature of a stochastic component in the acceleration process. Very recently, the log-parabolic law has been also applied to describe the spectral distribution and evolution of some Gamma-ray bursts \\citep{Massaro2011}.   The principal aim of the present paper is to investigate this scenario by extracting information on the acceleration processes using the curvature parameter measured in the observed synchrotron and Inverse  Compton (IC) spectra of Synchro-Self Compton (SSC) sources. We study the conditions in which the energy distributions of electrons, resulting from stochastic acceleration, can be approximated by a log-parabolic law and how its curvature evolves during their acceleration, and the role of IC cooling. We compare predictions from our theoretical descriptions with the curved spectra of some HBL objects.   In Sec. \\ref{sec:Logpar} we give an intuitive picture to take into account the effect of random fluctuations in the energy gain of particles and the role these play in determining the spectral curvature, as a consequence of the multiplicative central limit theorem,  and compare these results with the analytical solution of the diffusion equation, in the ``hard-spheres\" approximation. In Sec. \\ref{sec:MC} and Sec. \\ref{sec:numeric} we give a more physical description of the problem, using first a Monte Carlo approach, and secondly by solving numerically the momentum diffusion equation. We discuss the evolution of the curvature in the electron distribution as a result of momentum-diffusion before equilibrium is reached, and the role that synchrotron and IC cooling processes play on reaching the equilibrium. In Sec. \\ref{sec:ssc_evolution},  we study the peak energy, fluxes, and curvature, trends in the SED of both the synchrotron and IC emission, looking for the fingerprints of the stochastic component. In Sec. \\ref{sec:fit_data}, we show how our results can reproduce the spectral trends observed in a some HBLs, in particular we investigate the relation between the peak energy and the curvature, and between the peak energy and the peak flux. The good agreement between predictions and observed trends, confirms that the stochastic acceleration mechanism can play an important role in the physics of the blazars' jets and other SSC sources. ",
        "conclusions": "Broad band observations of non-thermal sources have shown that the spectral curvature at the peaks of their SEDs can now be measured with  good accuracy. In this paper, we have presented, using different approaches, the relevance of these data for the understanding of the competition between statistical acceleration and radiation losses. First, using a simple statistical approach and Monte Carlo calculations, we have shown that the log-parabolic energy distribution of the relativistic electron is a good picture in the first phases before  equilibrium is reached. In this case the curvature decreases with time and, therefore, for increasing peak energies. This evolution is confirmed by numerical solutions of the diffusion equation taking properly into account both stochastich acceleration and radiative SSC cooling. The major results can be summarised as follows.  The evolution of the electron energy distributions (Sec. \\ref{sec:numeric}) shows that: \\begin{itemize} \\item in the case of synchrotron and SSC cooling, and for all the values of $B$ and $R$, as long as the distribution is far from equilibrium, the trend on $r$ is dictated by $D_p$, and is well described by Eq. \\ref{eq:analyt_sol_curv}; \\item when the distributions approach  equilibrium, the value of $r$ is determined by the shape of the equilibrium distribution, which is a relativistic Maxwellian, with the sharpness of the cuf-off determined by both $q$ and the IC cooling regime; \\item in the case of $q=2$, and for equilibrium energies implying that  IC cooling happens either in the TH regime or in the extreme KN regime (IC cooling negligible compared to the synchrotron one), the numerical solution of the diffusion equation follows the analytical prediction ($f=1$, that holds for any $\\dot\\gamma\\propto\\gamma^2$), and the corresponding equilibrium curvature is $r_{3p}\\approx 6.0$ ($b_s\\approx 1.2$). In the case of $q=3/2$ the equilibrium curvature is $r_{3p}\\approx 3.0$  ($b_s\\approx 0.6$). These limiting values could be a useful observational test to find cooling dominated flares with the distribution approaching to the equilibrium; \\item when  cooling is in the intermediate regime between TH and KN and for the $q=2$ case, the condition $f=1$ fails, and the end values of $r$ decrease, strongly depending on the balance between $U_B$ and the seed IC photon energy ($U_{ph}$); numerical computations are necessary to evaluate the right value of $r$ at equilibrium. \\end{itemize}  The analysis of the spectral evolution of SSC emission (Sec.\\ref{sec:ssc_evolution}) shows that: \\begin{itemize} \\item changes of $D_{p0}$ (or $q$) imply that the curvature and peak energy of the synchrotron emission are anticorrelated; the $E_s$-$b_s$ trend can be phenomenologically described by Eq. \\ref{eq:Ep_b}; \\item The $E_c$-$b_c$ trend presents a clear signature of the transition from the TH to the KN regime. In particular when the IC scattering approaches the KN regime we observe a sharp change in the $b_c$, with a positive correlation with $E_c$, whilst in the TH regime the correlation is negative as in the case of the $E_s$-$b_c$; \\item the magnetic field plays a relevant role on the cooling process, and $B$ driven variations present relevant differences compared to those due to $D_{p0}$ (and $q$). \\end{itemize}  In particular, for what concerns the B driven case, we note first that the $E_s$-$S_s$ correlation follows the prediction of the synchrotron theory and shows the power-law relationship with $E_s\\propto(S_s)^{\\sim2.0}$. On the contrary, in the case of $D_{p0}$ and $q$ changes, we find $E_s\\propto (S_s)^{0.6}$. Another relevant difference in the $B$ driven case is the evolution of $S_c$. For the case of $D_{p0}$ and $q$ driven trends $S_c$ relates to $E_c$ through a power-law with exponent of about [0.7-0.8]. On the contrary, for the $B$ driven case with IC scattering in the full KN regime, the value of $E_c$ is almost constant and uncorrelated with $S_c$ (see Fig. \\ref{fig:Ep_vs_b_B_trend}), due to the kinematic limit of the KN regime. $E_c$ starts to decrease when $B$ is enough large to make dominant the cooling process. This is an interesting signature that could be easily checked in the observed data.    The comparison of the $E_s$-$b_s$ and $E_s$-$S_s$ trends, obtained through several X-ray observations of six HBL objects spanning a period of many years, with those predicted by the stochastic acceleration model, shows very good agreement. We are able to reproduce these long-term behaviours, by changing the value of only one parameter ($D_{p0}$ or $q$). Interestingly, the $E_s$-$S_s$ relation follows naturally from that between $E_s$ and $b_s$. This result is quite robust and hints at a common accelerative scenario acting in the jets of HBLs.  As a last remark, we note that very recently \\cite{Massaro2011} find also in the case of GRBs a $E_s$-$b_s$ trend, similar to that observed in the case of HBL objects. They measured values of the curvature up to 1.0, typically higher than in HBLs. It's interesting to note that the value of 1.0 is close to the limit of $\\sim$1.2, that we predict in the case of distributions approaching the equilibrium in either TH or KN regime, for $q=2$. \\\\"
    }
}